{
    "1402/1402.6308_arXiv.txt": {
        "abstract": "We present a {\\it Spitzer} MIPS study of the decay of debris disk excesses at 24 and 70 $\\mu$m for 255 stars of types F4 - K2. We have used multiple tests, including consistency between chromospheric and X-ray activity and placement on the HR diagram, to assign accurate stellar ages. Within this spectral type range, at 24 $\\mu$m, $13.6 \\pm 2.8 \\%$ of the stars younger than 5 Gyr have excesses at the 3$\\sigma$ level or more, while none of the older stars do, confirming previous work. At 70 $\\mu$m, $22.5 \\pm 3.6\\%$ of the younger stars have excesses at $ \\ge$ 3 $\\sigma$ significance, while only $4.7^{+3.7}_{-2.2}$\\% of the older stars do. To characterize the far infrared behavior of debris disks more robustly, we double the sample by including stars from the DEBRIS and DUNES surveys. For the F4 - K4  stars in this combined sample, there is only a weak (statistically not significant) trend in the incidence of far infrared excess with spectral type (detected fractions of 21.9$^{+4.8}_{-4.3}\\%$, late F; 16.5$^{+3.9}_{-3.3}\\%$, G; and 16.9$^{+6.3}_{-5.0}\\%$, early K).   Taking this spectral type range together, there is a significant decline between 3 and 4.5 Gyr in the incidence of excesses with fractional luminosities just under $10^{-5}$. There is an indication that the timescale for decay of infrared excesses varies roughly inversely with the fractional lumnosity. This behavior is consistent with theoretical expectations for passive evolution. However, more excesses are detected around the oldest stars than is expected from passive evolution, suggesting that there is late-phase dynamical activity around these stars. ",
        "introduction": "Understanding planetary system formation and evolution is one of the major initiatives in astronomy. Stars form surrounded by protoplanetary disks of primordial gas and dust where planets grow. The material in these disks that does not fall into the star either collects into planets or is dissipated by processes such as photo-evaporation \\citep[e.g.,][]{clarke2001} and tidal forces from planets \\citep[e.g.,][]{bryden1999}, typically in less than 10 Myr \\citep{williams2011}. The evolution of the systems is not complete, as shown by the events that led to the formation of our Moon and to the Late Heavy Bombardment \\citep{tera1974}, long after the protoplanetary disk cleared from the Sun. It is very difficult to observe these later stages of evolution directly. However, after dissipation of the protoplanetary disks, a relatively low level of dust production can be sustained through debris produced in planetesimal collisions  to form planetary debris disks \\citep{wyatt2008}, which can be detected readily in the infrared over the entire range of stellar lifetime (to 10 Gyr). Debris disks are our best current means of studying planet system evolution over its entire duration.  About 20\\% of the nearby stars harbor debris disks above current detection limits \\citep{habing2001,trilling2008,carpenter2009,gaspar2013,eiroa2013}. The likelihood of a detectable debris disk at 24 $\\mu$m depends on age, with a higher percentage around young stars than older ones \\citep[e.g.,][]{habing1999,rieke2005,gaspar2013}. The expected timescale for evolution of the disk components that dominate in the far infrared indicates that, depending on disk location and density, there might also be a detectable decay in disk incidence at 70 $\\mu$m \\citep{wyatt2008}. However, it has proven difficult to confirm this prediction definitively \\citep{wyatt2008, bryden2006, trilling2008,carpenter2009}. \\citet{gaspar2013} have recently demonstrated a drop in infrared excess on the basis of comparison with the predictions of a theoretical model. If such a decay can be confirmed by an alternative analysis, it would help substantially to constrain models and particle properties of debris disks.  In this paper, we explore debris disk evolution in the far infrared (70 - 100 $\\mu$m) using a large sample of stars to allow reaching statistically robust conclusions. The paper presents high-quality homogeneous data reductions for a large sample of stars observed with {\\it Spitzer}, analyzes their behavior, and then combines this sample with the {\\it Herschel}-observed DEBRIS \\citep{matthews2010} and DUNES \\citep{eiroa2013} samples. We take care in determining stellar ages, since the apparent rate of decay can be strongly influenced by mis-classification of the ages (a small number of young stars mistakenly identified as old ones might dominate the detections among the \"old\" sample). Section 2 describes the selection of the {\\it Spitzer} sample, the determination of stellar ages, and the reduction of the infrared measurements. We determined debris-disk-emitted excesses as described in Section 3. Section 4 presents the analysis of the behavior of these excesses with stellar age, and Section 5 merges this sample and additional {\\it Spitzer} data with {\\it Herschel} observations to study debris disk evolution in a total sample of about 470 stars. Our conclusions are summarized in Section 6. ",
        "conclusions": "We present a sample of 238 F4 - K2 stars with high-quality {\\it Spitzer} observations through 70 $\\mu$m, and for which we have been able to determine accurate ages from the chromospheric activity index, X-ray emission, and placement on an HR-diagram. Consistent with previous work, we find that the incidence of excesses at 24 $\\mu$m decays dramatically from $13.6 \\pm 2.8$ \\% for stars $<$ 5 Gyr in age to none for older stars. At 70 $\\mu$m, we also find a significant reduction in the incidence of excesses, from 38, or $22.5 \\pm 3.6$ \\% to 4, or $4.7^{+3.7}_{-2.2}$ \\%. We have evaluated the significance of this apparent reduction in the incidence of 70 $\\mu$m excesses in two ways, both of which correctly account for the variable detection limits and resulting censoring of the data. The first is the traditional Kaplan-Meier test, which  indicates a decay by greater than 2$\\sigma$ significance, nominally by a factor of 4.5. In the second method, we determined the distribution of excess ratios among the young (ages $<$ 5 Gyr) stars and assumed it could be treated as a probability distribution to determine if a specific star is likely to have an excess above its individual detection limit. This model was used to predict the number of excesses that would have been detected in the old ($>$ 5 Gyr) sample if they behaved identically to the young sample. This approach also predicts a decay at greater than 2 $\\sigma$ significance, and nominally by a factor of five. The two methods therefore agree well on the existence and amount of this decay. Within this sample, there are two stars with excesses at 24 $\\mu$m but no strong excess at 70 $\\mu$m, HD 1466 and HD 13246.  We then combined the {\\it Spitzer}-only stars with the measurements for the DEBRIS and DUNES programs (for the former using both {\\it Spitzer} and {\\it Herschel} data to improve the weight of the measurements). This combined sample is doubled in size, allowing us to reach robust conclusions about debris disk evolution in the far infrared. Within this larger sample, we examine the relative incidence of excesses around stars of types F4 - K4, finding no statistically significant trend with spectral type. We therefore analyze the evolution of the entire sample together. We find that, at fractional debris disk luminosity of just under $10^{-5}$, the incidence of far infrared excesses declines significantly between 3000 and 5000 Myr, while for fractional luminosity of just under $10^{-4}$, the decline is at an order of magnitude younger ages. This behavior is consistent with a theoretical model for passive disk evolution. However, there appear to be more excesses for the oldest stars than predicted by passive-evolution models. These excesses are distributed over the full range of spectral types in our sample. This behavior suggests that the oldest debris disks are activated by late-phase dynamical activity."
    },
    "1402/1402.0697_arXiv.txt": {
        "abstract": "{Halo masses and concentrations have been studied extensively, by means of N--body simulations as well as observationally, during the last decade. Nevertheless, the exact form of the mass--concentration relation is still widely debated. One of the most promising method to estimate masses and concentrations relies on gravitational lensing from massive halos. Here we investigate the impact of the mass--concentration relation on halo peak abundance in weak lensing shear maps relying on the aperture mass method for peak detections. After providing a prescription to take into account the concentration dispersion (always neglected in previous works) in peak number counts predictions, we assess their power to constrain the mass--concentration relation by means of Fisher matrix technique. We find that, when combined with different cosmological probes, peak statistics information from near--future weak lensing surveys provides an interesting and complementary alternative method to lessen the long standing controversy about the  mass-concentration relation. } ",
        "introduction": "Weak gravitational lensing (WL) has emerged over the last decade as one of the most promising methods for testing cosmology and gravity, and unveiling the nature of Dark Energy (DE) and Dark Matter (DM). It relies on the accurate measurement of the small shape distortions of background galaxies due to the bending of light by intervening matter distribution. The smallness of these distortions requires high-quality data  and a statistical approach so that the lensing shear signal needs to be measured on a large number of sources.  A number of surveys (e.g. GaBoDS \\cite{Hett07}, CFHTLS\\footnote{http://www.cfht.hawaii.edu/Science/CFHLS}) have already shown the WL capacity for constraining cosmological models through cosmic shear measurements (see e.g. \\cite{Hoekstra06,Semboloni06,Hett07}). Several other, more ambitious, surveys both ground-- (KIDS\\footnote{http://kids.strw.leidenuniv.nl}, PanSTARRS\\footnote{http://pan-starrs.ifa.hawaii.edu}, DES\\footnote{http://www.darkenergysurvey.org}, LSST\\footnote{http://www.lsst.org}) and space-based (Euclid\\footnote{http://www.euclid-ec.org}, WFIRST\\footnote{http://wfirst.gsfc.nasa.gov}) are being performed or planned. They will map hundreds of millions of galaxy redshifts and billions of galaxy images. For instance, the Euclid mission \\cite{euclid1,euclid2} aims to map half of the sky in imaging as well as in broad-band spectroscopy up to a redshift $z \\sim 2-3$ with a median redshift of the order of unity. The information contained in these data will permit to measure the matter clustering with unprecedented accuracy.  Among the WL probes, it has been shown that the abundance of peaks in shear maps is a sensitive probe of cosmology \\cite{kruse99}. This was tested by \\cite{Reblinsky99} and \\cite{Dietrich10} using ray--tracing in N--body simulations. Shear peaks are regions with high signal--to--noise ratios, $snr$, associated to massive halos (or galaxy clusters) or produced by the alignment of smaller mass concentrations along the line--of--sight. Their number and spatial distribution, therefore, carry information about the underlying cosmology through fundamental parameters as the total matter density of the universe $\\Omega_m$, the normalization of the power spectrum $\\sigma_8$, and the evolution of the DE equation of state $w(z)$.  Unlike other available techniques for detecting  massive halos (e.g. optical and X--ray identification, Sunyaev-Zeldovich effect, etc), WL does not require any assumption on their dynamical or evolutionary state. Its advantage is to be only sensitive to the mass along the line--of--sight offering, in principle, the opportunity to construct mass--selected halo samples, which can be compared directly with theoretical predictions (e.g. N--body simulations) without assuming any mass--observable relation. Nevertheless, that WL provides truly mass--selected halo samples it is not strictly true. Large scale structures (LSS), into which massive halos are embedded, contribute to the lensing signal. Usually, these  projection effects add noise to the halo signal or can result in false positive peaks in the shear map \\cite{Hoekstra01,Hoekstra03,Hoekstra11}. Further, intrinsic ellipticities of galaxies introduce an irremovable noise. All these sources of noise then compromise purity and completeness of halo samples selected by WL.  Various theoretical aspects of WL halo detection have been widely investigated by a number of authors \\cite{hamana04,Dietrich10,marian06,maturi10,maturi11,cardone}. The search for shear peaks in data is exemplified in some works \\cite{hett05,dahle06,wittman06,gavazzi07,schirmer07,miyazaki07,berge08,abate09,shan11} although, on the observational side, halo lensing studies have been mainly focused on mass determination and halo properties, e.g. density profile and concentration \\cite{okabe10,israel10}.  A basic method for shear peak detection, to which we will refer in the following, is provided by the aperture mass ($M_{ap}$) technique proposed by \\cite{S96}. $M_{ap}$ is a measure of the lensing signal in a shear map smoothed with a suitable filter. The method then relies on the usual approach of searching for points of local maximum (peaks), with $snr$ higher than a given threshold, in the smoothed map.  Here, we investigate the impact of the halo concentration on $M_{ap}$ and shear peak counts, as compared to that of some cosmological parameters (i.e. $\\Omega_m$ and $\\sigma_8$). As already noted by \\cite{king11} and detailed in the following, the shear peak abundance, is quite sensitive to halo concentration and its relation with the halo mass. This seems to point out  peak counts as a promising and complementary method to delineate the, still debated, relation between halo mass and concentration ($M-c$ relation). Our main interest here is to further deepen this point quantifying the power of peak function to constrain the $M-c$ relation.  Halo concentrations have been studied extensively, by means of N--body simulations as well as observationally, during the last decade. The general trend is that the mean halo concentration shows a strong correlation with the halo mass. Typically, it declines with increasing the mass and redshift in accordance with the idea that the central density of halos reflects the mean density of the universe at the time of their formation. Therefore, halos collapsing earlier are expected to be denser than the more massive halos collapsing later.  However, there is enough uncertainty on the exact form of the $M-c$ relation. Beside discrepancies between theoretical predictions and observations, differences are found even when comparing simulation results of different groups as well as different sets of observations \\cite{B01,eke01,Comerford07,neto07,gao08,duffy08,maccio08,mandelbaum08,oguri09,okabe10,klypin11,prada11,meneghetti13}.  One of the main predictions of N--body simulations is that the density profile of halos assembled hierarchically and close to virial equilibrium, can be well approximated by a Navarro, Frenk \\& White (NFW) profile \\cite{NFW}, regardless of halo mass and the details of the cosmological model, all the cosmological informations being contained in correlation between the parameters of the NFW profile, i.e. in the $M-c$ relation. The median $M-c$ relation predicted by simulations is well described by a power--law \\cite{dolag04,gao08,duffy08,maccio08,zhao09}.  However, the variety of individual halo aggregation histories causes a scatter in concentration which can be modeled by a log--normal distribution with a variance ranging from $\\sim 0.15$ up to $\\sim 0.30$ depending on the degree of relaxation of the halo \\cite{Jing00}. On the other hand, the redshift evolution of the concentration is less established. Recent results seem to favor a weaker redshift dependence than previously supposed \\cite{duffy08,maccio08,MC11}.  Probably, the best current results from simulations are those by \\cite{klypin11} and \\cite{prada11}. They found larger concentrations than those reported in \\cite{maccio08} and \\cite{duffy08}, with a reasonable agreement for galaxy--size halos ($\\sim 10-15$\\% difference) but substantially larger values ($\\sim 40-50$\\%) for cluster--size halos, more compatible with recent X--ray and kinematic observations. Furthermore, these high--resolution large-volume simulations ({\\it Bolshoi} and {\\it MultiDark}, see http://www.mutidark.org) seem to indicate a $M-c$ relation more complex than previously conceived showing a novel feature: at high redshift, concentration first declines with increasing mass, then flattens and  increases slightly at higher masses.  On the observational side, many different methods have been used to study masses and concentrations resulting in an even more controversial picture and making the comparison between observations and theory somewhat ambiguous.  Gravitational lensing is one of the most promising methods to estimate mass and concentrations. Strong and weak lensing have been used individually or in combination. More classical methods relies on kinematic tracers, e.g. galaxies (\\cite{rines06,WL10} and references therein) and the hydrostatic analysis of X--ray intensity profile of halos. However, all these methods are susceptible to bias effects and are presumably affected by systematics not fully understood and/or modeled, yielding discrepant results.  For example, \\cite{mandelbaum08} used the stacked weak lensing signal from galaxies, groups and clusters in the Sloan Digital Sky Survey (SDSS) finding a $M-c$ relation slope consistent with the simulations but a $2\\sigma$ lower normalization. \\cite{oguri09} performed a combined weak and strong lensing analysis of a sample of clusters, reporting concentrations with a $7\\sigma$ excess above the simulation predictions. \\cite{okabe10} performed a weak lensing analysis of X--ray selected clusters reporting a slope somewhat steeper than in simulations although with an error of $\\sim 50$\\%. Other studies based on lensing and X--ray data can be found in \\cite{Comerford07,buote07,SA07,johnston07,broadhurst08,vikhlinin09,ettori11}.  Despite the qualitative agreement with simulation predictions, the general picture which emerges from observations can be summarized briefly as follows. In almost all cases, the observed slope of the $M-c$ relation is consistent with or steeper than theoretical predictions. Moreover, strong lensing measures of massive clusters and X--ray analysis give a normalization factor higher than predicted by simulations, strong lensing concentrations being systematically larger than X--ray concentrations. On the other hand, weak lensing methods seem to point out a normalization lower than that found in simulations.  For what concerns the lensing measurements, the origin of these discrepancies should be searched in some orientation and shape biases \\cite{oguri05,sereno11,giocoli}. It has been shown that neglecting halo triaxiality can lead to over-- and under--estimates of a factor of $2$ in concentrations and up to $50$\\% in halo mass other than underestimation of statistical uncertainties \\cite{corless09}. In addition, as already outlined above, projection of structures along the line--of--sight can result in apparently high concentration (e.g. \\cite{KC07}).  Further complications are due to baryonic processes. Although baryon physics is not expected to drive the structure formation process on very large scales, its effects are likely to be important for low--mass halos or in the central region of larger objects. It has been shown that baryonic feedback and cooling can alter halo profiles. In particular, baryonic cooling (see \\cite{gnedin11} and reference therein) can be responsible for the excess in concentration observed in groups and low-mass clusters \\cite{sereno10,fedeli11}.  Dynamical and X--ray techniques assume the halos are in virial and/or hydrostatic equilibrium. Most of real halos are not relaxed and have complex substructures making measurements difficult to interpret. Mergers, active galactic nuclei, cosmic rays, magnetic fields, turbulence and bulk motion of the gas can compromise the hydrostatic equilibrium leading to underestimates of X--ray masses \\cite{evrard96,dolag05,rasia06,nagai07,rasia13}.  Although recent hydrodynamical simulations include a number of these non--gravitational processes, they are very difficult to model. Furthermore, many of them take place at scales too small to be resolved by simulations demanding radical simplifications and approximations.  It is also worth noticing that the above results were derived in the contest of the standard cosmological model ($\\Lambda$CDM). Discrepancies can also arise if the data--model comparison has been made for the wrong cosmology. A number of studies have shown that alternative DE cosmologies or modified gravity models can lead to a different evolution of matter perturbations affecting the halo concentrations. For example, scalar field DE models with an earlier structure formation resulting in more concentrated halos have been considered in \\cite{dolag04,io1,io2,baldi,baldi2,boni}.  The paper is organized as follows: the $M_{ap}$ method for shear peak detection is reviewed in Section \\ref{Map} while in Section \\ref{PF} we investigate the impact of the $M-c$ relation on the probability distribution function of $M_{ap}$ and provide a prescription to analytically calculate it taking into account the dispersion in $M-c$ relation. The impact on peak number counts is then investigated and compared to that of $\\Omega_m$ and $\\sigma_8$. In Section \\ref{fisher} we perform a Fisher matrix analysis in order to assess the capability of upcoming weak lensing survey in constraining the $M-c$ relation by means of peak counts. Section \\ref{concl} is devoted to discussion and conclusions. ",
        "conclusions": "\\label{concl}  Halo masses and concentrations have been studied extensively, by means of N--body simulations as well as observationally, during the last decade. Nevertheless, the exact form of the $M-c$ relation is still widely debated. One of the most promising method to estimate masses and concentrations relies on gravitational lensing, in particular shear peak counts from near--future WL surveys seem a promising and complementary method to delineate the features of the $M-c$ relation, as pointed out in this work.  The method for shear peak detection we have considered here relies on the $M_{ap}$ statistics. Hence, we have firstly investigated the impact of the $M-c$ relation on the PDF of $M_{ap}$. In particular, we have provided a prescription to properly take into account the $M-c$ relation dispersion, always disregarded in previous work, in the PDF. We have then shown that such scatter can affect significantly the PDF and, consequently, the $snr$ for peak detection.  Secondly, we have investigated the impact of the $M-c$ relation on peaks number counts as compared to that of the cosmological parameters $\\sigma_8$ and $\\Omega_m$.  Finally, we have performed a Fisher matrix analysis in order to assess the capability of an Euclid--like survey in constraining the $M-c$ relation and cosmological parameters by using peak number counts. We have found that: i) peak function alone provides constraints on $\\sigma_8$ and $\\Omega_m$ which are competitive with those obtained from different probes. On the other hand, only the normalization $c_0$ of the $M-c$ relation is reasonably (but not so strongly) constrained if peaks with $snr>4$ are considered; ii) adding prior informations on $\\sigma_8$ and $\\Omega_m$ as inferred from current/near--future data, constraints on $M-c$ relation improve significantly so that peak statistics seems helpful in discriminating among the wide variety of $M-c$ relation fits found in literature although no substantial new information on the slope $\\alpha$ is inferred.  Although results are quite encouraging, it is worth stressing that the Fisher matrix analysis here presented disregards the effects of the halo sample variance on the error estimates, only Poisson noise being considered. Nevertheless, sample variance might be significant for low mass halos at low redshifts where their abundance is much higher thus reducing the Poisson noise. To yield more accurate predictions, the covariance matrix for number counts $C_{ij}$ \\footnote{it should not be confused with the parameter covariance matrix of eq. (\\ref{sigpiv}) for which we use the same simbol.}, entering the FM expression, should then be written as: $C_{ij} = P_{ij}+S_{ij}$ where $P_{ij}=\\delta_{ij}N_i^{fid}$ and $S_{ij}=N_i^{fid}N_j^{fid}b_ib_j \\sigma_{ij}^{DM}$ are the Poissonian and sample variance contributions respectively (here $N_i^{fid}$ and $b_i$ are the expected number of peaks and the mean halo bias in the $i$--th redshift bin, while $\\sigma_{ij}^{DM}$ is the covariance of the dark matter fluctuations between the bins $i$ and $j$, se e.g. \\cite{hukra,takada,valageas} for more details).  Sample variance effects on halo number counts have been considered in previous works, e.g. \\cite{hukra,takada,valageas,lima,masamune}. In particular,  \\cite{hukra,valageas} conclude that sample variance is generally comparable or greater than Poisson noise for number counts above mass threshold $M_{th} \\lesssim 1-4 \\cdot 10^{14} h^{-1} M_\\odot$ depending on the survey width and depth. In our case, the mass thresholds corresponding to given $snr^*$ and $z$ can be deduced from the middle panel of fig. \\ref{snzM}. According to  \\cite{hukra,valageas}, for $snr^* \\geq 4$ we expect sample variance to have only moderate effects in the redshift range $\\Delta z=0.15-0.45$ where the counts signal is maximum, becoming negligible at the increasing of $z$. For a raw estimate, we can assume $C_{ij}$ to be diagonal (i.e. halo correlation length smaller than the bin width) and $S_{ij}=P_{ij}$ ($S_{ij}=2P_{ij}$) independently of the redshift. This yields a worsening of the parameter constraints by a factor $\\sqrt 2$ ($\\sqrt 3$).  However,  \\cite{hukra,takada,valageas} neglect the information arising from the cosmological dependence of the sample variance. When this information is included and considered as ``signal'', it could cause a smaller degradation in the constraints or, in case, even improvements \\cite{lima,masamune}. In this case, the full FM can be well approximated as:  $$ F_{ij} = {\\bf N}^T_{,i} {\\bf C}^{-1} {\\bf N}_{,j} + \\frac{1}{2} {\\textrm {Tr}} \\left[{\\bf C}^{-1}{\\bf S}_{,i}{\\bf C}^{-1}{\\bf S}_{,j} \\right]+ \\frac{\\delta_{ij}}{\\sigma_p^2(p_i)} $$ where ${\\bf N}=(N^{fid}_1,...,N^{fid}_{n_{bin}})$ and $_{,i}$ denotes the derivative with respect to the parameter $p_i$. It is worth stressing that the second term, neglected in \\cite{hukra,takada,valageas}, beside the information about the cosmological dependence of the sample variance, also contains information about its dependence on $M-c$ relation trough the derivative of $N_i^{fid}$ and possibly of $b_i$. Indeed, it was pointed out in \\cite{wechsler,jingsu} that halo bias could depend not only on the halo mass but also on additional halo properties, among which the concentration.  Adding sample variance information could then, in principle, provide an opportunity to improve measurements of halo concentration (and cosmological parameters) with WL peak statistics through self--calibration techniques for $M-c$ relation similar to those for mass--observable relation (see e.g. \\cite{lima,masamune}). This intriguing point, which deserve a detailed analysis, is left for future works.     A second issue concerns the assumption that detected peaks can be properly assigned to redshift bins. The redshift information, however, can not be extracted from WL data only, but additional data need to be included. In order to split the peak sample in redshift bins, optical finders (see e.g. \\cite{postman,koester,milke,bella}) could serve the purpose since they return reasonable redshift estimates, especially because they can be applied on the same optical data retrieved for lensing. Although the correlation between optical detections and WL peaks is not trivial, in a Euclid--like survey, one can look at the position of the peaks in the optical images, and then deduce the peak redshift $z$ from the redshifts of the galaxies closest to the peak position. In a first approximation, we can suppose that this method provides a Gaussian probability distribution function for $z$ with negligible bias and variance $\\sigma_z = \\sigma_0(1+z)$, and then assume a peak to be correctly assigned to a redshift bin roughly asking that the $3-\\sigma$ uncertainty on $z$ is smaller than the bin width.  For our assumed value $\\Delta z = 0.1$, this translates in $\\sigma_0 \\leq 0.03$, a precision which could be likely achieved if $z$ is spectroscopically measured ($\\sigma_0 \\simeq 0.001$ according to the Euclid red book), but could be too demanding if one relies on photometric redshift methods ($\\sigma_0 \\simeq 0.05$ for Euclid). In this second case, one should add a non Poissonian uncertainty on $N_{halo}(snr > snr^*, z, \\Delta z)$ in the above Fisher matrix analysis and resort to the pull statistics \\cite{campanelli}. On the other hand, it is likely that the precision of the inferred redshift also depends on the $snr$ of peak so that the net effect should be included in the analysis by convolving the theoretically computed $N_{halo}$ with an empirically determined selection function.  A further issue concerns the assumption of NFW spherical halos. Real halos are not spherical and not all relaxed so that deviations from the universal NFW profile are expected. The effects of the diversity of dark matter distributions in individual halos on peak counts has been investigated in \\cite{gruen,hamana12} by means of numerical simulations and analytic methods. The noise originated from halo shape is found to be comparable to the statistical noise discussed in section \\ref{pdf}. Furthermore, halo orientations cause a systematic bias in the peak heights. In order to include these effects in the theoretical predictions, one should consider a triaxial halo model (see \\cite{jingsuto}) which is shown to work quite well \\cite{hamana12}.  It is also worth mentioning that, in the FM analysis, we have only investigated a subset of the full parameter space. Here, we were mainly interested in how peak statistics constrain the $M-c$ relation so that those parameters, as $H_0$ and $n_s$, which have a minor impact on peak number counts and are expected to be scarcely constrained, were held fixed to their fiducial values. Nevertheless, when dealing with real data, the full parameter space should be considered. Allowing for a large number of parameters to be varied introduces further degeneracies widening the confidence regions. This degradation in the constraining power, can then be compensated by complementing peak statistics with further datasets, e.g. SNeIa and CMB.  Investigating the above issue is however outside our aims here."
    },
    "1402/1402.2976_arXiv.txt": {
        "abstract": "The massive exploitation of cosmic voids for precision cosmology in the upcoming dark energy experiments, requires a robust understanding of their internal structure, particularly of their density profile. We show that the void density profile is insensitive to the void radius both in a catalogue of observed voids and in voids from a large cosmological simulation. However, the observed and simulated voids display remarkably different profile shapes, with the former having much steeper profiles than the latter. Sparsity can not be the main reason for this discrepancy, as we demonstrate that the profile can be recovered with reasonable accuracy even with very sparse samples of tracers. On the other hand,  the observed profile shows a significant dependence on the galaxy sample used to trace the matter distribution. Samples including low-mass galaxies lead to shallower profiles with respect to the samples where only massive galaxies are used, as faint galaxies live closer to the void centre.  We argue that galaxies are biased tracers when used to probe the matter distribution within voids. ",
        "introduction": "Large redshift surveys \\citep{York00, Colless01}  and cosmological simulations \\citep{Bond96} have revealed that galaxies are distributed inside a cosmic web of walls, filaments and compact clusters. Such a web encloses large underdense regions, referred to as cosmic voids.  Voids were first recognized in the earliest redshift surveys \\citep{Gregory78, Kirshner81}  as huge empty holes in the galaxy distribution. Nowadays, there is a general consensus in that voids occupy most of the volume of the Universe \\citep{Sheth04, Vandew11, Pan12}, although there is  not yet an agreement on how a genuine void should be defined. Several void finders, which are  based on different principles, have been developed. Voids can be identified as spherical regions devoid of galaxies/haloes  \\citep{Gottlober03, Patiri06b, Varela12} or underdense regions, relying on the continuous density field \\citep{Plionis02, Colberg05}. More complex algorithms  able to capture the complex morphology of voids also exist \\citep{Platen07, Neyrinck08, AragonCalvo13}.  Despite their different definition of voids, all these void finders agree in that voids ere extremely empty in the centre and show a sharp increase in the density towards the voids edges (e.g. \\citealt{Colberg08}).   Voids are believed to originate from negative density fluctuations in the primordial density field. As a result of their underdensity, they are subject to an effective repulsive peculiar gravity, causing their expansion. As a consequence of such an expansion, the matter within the voids evacuates from the interior and accumulates to the boundaries. This leads to void density profiles that evolve towards a reverse top-hat shape \\citep{Sheth04}.  A considerable appeal of cosmic voids is their potential in probing cosmological parameters. In particular, being almost devoid of matter, they are extremely sensitive to the nature of dark energy.  Indeed, the void ellipticity and its evolution through cosmic time are intimately connected with the local tidal tensor, which, in turn, depends on the dark energy content  \\citep{Park07, Lavaux10, Bos12}. Voids  are also the ideal candidate for probing the expansion history of the Universe through the Alcock-Paczynski test \\citep{AP79}, using the average shape of stacked voids \\citep{Lavaux12, Sutter12a}. The application of such a test to the voids that will be identified in the future Euclid survey \\citep{Laureijs11} promises to outperform Baryonic Acoustic Oscillation by an order of magnitude in accuracy.  The huge potentiality of voids for precision cosmology requires a robust knowledge of their internal structure, particularly of the density profiles. Works based on cosmological simulations  \\citep{Colberg05, Ricciardelli13} indicate that the void density profile is universal. As such,  it does not depend on void size. On the observational side, the ideal approach to directly constrain the void density profile is through the  weak lensing signal of stacked voids \\citep{Krause13}. However, the number of voids available from spectroscopic catalogues is still limited to provide a robust measurement of the signal \\citep{Melchior13}. At present, we can only rely on the galaxy distribution to trace the density within voids \\citep{Sutter12b}. Thus, to robustly assess a void model to describe the universal density profile, one also needs to   assess the systematic effects arising from the use of the sparse galaxy sampling.  In a previous work (\\citealt{Ricciardelli13}, hereafter RQP13), we have shown, by means of a cosmological simulation, that a two parameters law can be used to fit the density profile of voids of any size, density, morphology and redshift.  The best-fit parameters show some dependence on redshift, density, and, on a less degree, morphology, but they are almost independent on the void size, although the limited statistics prevented us to draw robust conclusions.  In this work, we want to test this model and its dependence on void radius, against an observed catalogue of voids and a larger simulation, thus dramatically increasing the statistics. In doing so, we provide a robust determination of the systematic effects arising when using the sparse distribution of void galaxies as  density tracers.  The structure of the paper is as follows. In Section \\ref{sdss} we introduce our catalogue of observed voids, in Section \\ref{sim} we describe the simulation used and our void identification procedure. The results on the void density profiles are discussed in Section \\ref{profiles}. We conclude in Section \\ref{conclu}. ",
        "conclusions": "We have robustly assessed the universality of void density profiles, by means of a catalogue of observed voids and a large cosmological simulation.  The observed void catalogue has been drawn from the SDSS database, and includes spherical voids whose radius is larger that  $7 \\, h^{-1}\\,  Mpc$ \\citep{Varela12}. To measure the density profiles in these  voids, we rely on the luminous galaxies. As a matter of comparison, we have performed a large cosmological simulation with the code MASCLET, devoted to follow the formation and evolution of the low-density regions. This simulation has been designed to target, with sufficient statistics, voids spanning a wide range of radius. To this aim, we have simulated a large volume, having a comoving side length of $512 \\, h^{-1}\\,  Mpc$,  with only one level of refinement in the AMR grid, reaching the spatial resolution of $0.5 \\, h^{-1}\\,  Mpc$. Since this  modest resolution does not allow to follow the formation of structures in the simulated box, void galaxies in our simulation do not form. Therefore, we adopt as density tracer the continuous density field or, where a sparse distribution of tracers is needed, the dark matter particles within the void regions.  The void density profiles recovered by means of the observed and simulated voids share the same qualitative shape, showing a significant underdensity in the centre and a sharp density increase approaching the void edges. Both profiles can be well described by the functional form proposed in RQP13.  However, the observed profile is significantly steeper than the simulated one.  To figure out the reasons for the steepness of the observed profiles, we have assessed the impact of the number and type of tracers on the resulting  density profile. The sparsity of the density tracers has been investigated by means of subsamplings of the simulated voids, populated with an increasing number of particles. We have shown that even in the  less populated void samples, the original density profile can be recovered with reasonable accuracy. Stacks built with a limited number of voids or sparsely populated present a significant noise at small radii, but no systematic effect with the number of voids/tracers is observed. The low impact of the sparsity of the tracers on the internal void density profiles has been pointed out also by \\citet{Sutter13}, using both dark matter particles and haloes as density tracers.  Nevertheless, we observe that the profile shape can have a significant dependence on the type of galaxies used to trace the matter distribution. Within the observed voids, the density profiles recovered by means of faint samples of galaxies are shallower than those determined through the brighter galaxies. The reason for that lies in the galaxy mass segregation within voids. In fact,  faint galaxies are those living closer to the void centre and, thus, allow to probe the matter distribution even in the innermost part of the voids.  The strong impact of the type of galaxies chosen to trace the density, forces us to use an homogenous sample of galaxies and voids, limited in volume and magnitude, to assess the dependence of the void density profile on the void radius. With such a sample, we have demonstrated the insensitivity of the observed void profile on void radius. Likewise, by using our simulated sample of voids, we do not observe any dependence of the profile shape on the void size, and the same best-fit can correctly describe voids whose size ranges from 7 to $\\sim 50 \\, h^{-1}\\,  Mpc$.   Finally, we note that the difference in profile between the observed and simulated voids can not be driven by the different algorithms used to identify voids. Indeed, the density profile of our SDSS stack is very similar to the profile published in Pan et al. (2012), using the same SDSS DR7 dataset, albeit with a completely different void finder. Moreover, our simulated void density profiles are in remarkable agreement with the simulations of \\citet{Colberg05}, where voids are identified through spherical underdensities. Therefore, we argue that the difference between observed and simulated void density profiles is a robust result and is due to the biased tracers used, when relying on the observed galaxies. To corroborate this hypothesis, we definitely need high resolution simulations, capable to follow structure formation in the most rarefied regions of the Universe."
    },
    "1402/1402.1849_arXiv.txt": {
        "abstract": "Accretion disk winds are thought to produce many of the characteristic features seen in the spectra of active galactic nuclei (AGN) and quasi-stellar objects (QSOs). These outflows also represent a natural form of feedback between the central supermassive black hole and its host galaxy. The mechanism for driving this mass loss remains unknown, although radiation pressure mediated by spectral lines is a leading candidate. Here, we calculate the ionization state of, and emergent spectra for, the hydrodynamic simulation of a line-driven disk wind previously presented by \\cite{pk04}. To achieve this, we carry out a comprehensive Monte Carlo simulation of the radiative transfer through, and energy exchange within, the predicted outflow. We find that the wind is much more ionized than originally estimated. This is in part because it is much more difficult to shield any wind regions effectively when the outflow itself is allowed to reprocess and redirect ionizing photons. As a result, the calculated spectrum that would be observed from this particular outflow solution would not contain the ultraviolet spectral lines that are observed in many AGN/QSOs. Furthermore, the wind is so highly ionized that line-driving would not actually be efficient. This does not necessarily mean that line-driven winds are not viable. However, our work does illustrate that in order  to arrive at a self-consistent model of line-driven disk winds in  AGN/QSO, it will be critical to include a more detailed treatment of radiative transfer and ionization in the next generation of hydrodynamic simulations. ",
        "introduction": "\\label{introduction} \\begin{figure*} \\includegraphics{fig1_colour.eps} \\caption{The density (colours) and poloidal velocity (arrows) structure of the PK04 model. Radial lines delineate the three zones described in the text. } \\label{lin_den} \\end{figure*}  Outflows are found in a vast range of accreting astrophysical objects, from protostars to active galactic nuclei (AGN). In all of these settings, quantifying the mass and energy flows involved is key to understanding how the accreting objects evolve and interact with their local environment. In the case of AGN, the observed outflows can be split into two main classes: highly collimated, relativistic jets, and slower moving ($v \\ltappeq 0.2c$), but more massive `disk winds' driven from the surface of the accretion disk surrounding the central supermassive black hole. These disk winds have been proposed as the underlying structure responsible for many observed AGN spectral features, including the broad absorption lines seen in a significant proportion of QSOs \\citep[e.g.][]{knigge08}, the so called broad absorption line quasars (BALQSOs). However, the geometry of these winds, and even the mechanism by which they are launched, are still not known.  Several mechanisms have been proposed to produce the accelerating force for disk winds in AGN, including gas/thermal pressure \\citep[e.g.][]{weymann_82, begelman_91, begelman_83,krolik_kriss}, magnetocentrifugal forces \\citep[e.g.][]{blandford_payne_82, pelletier_pudritz} and radiation pressure acting on spectral lines (``line-driving'') \\citep[e.g.][]{shlosman_85,murray_95}. All of these mechanisms have been shown to produce outflows in suitable conditions. However, line driving is particularly attractive, since the requisite force in this case arises naturally when (primarily ultraviolet [UV]) photons produced by the central engine are scattered by strong resonance lines. Many such lines are observed directly in the UV spectra of (BAL)QSOs, so line driving must certainly be acting in some measure. Moreover, if line driving is the dominant acceleration mechanism, it can produce a unique signature in the profiles of absorption lines produced in the wind (the so-called ``ghost of $\\rm{Ly\\alpha}$'', e.g. \\citealt{arav_95,arav_96}), and this may already have been seen in several BALQSOs (\\citealt{north_knigge_goad}, but also see \\citealt{cottis_10}).  Despite these circumstantial reasons for favouring line-driving as the mechanism for the production of disk winds in AGN, there is one main challenge to the model. For line-driving to be efficient, the accelerated material has to be in a moderately low ionization state, despite its proximity to the intense X-ray source at the center of the accretion disk. If the gas becomes over-ionized, the ionic species that can most effectively tap into the momentum of the radiation field are simply not present.  \\cite{murray_95} proposed one solution, suggesting that `hitchhiking gas', material interior to the main outflow and itself accelerated by pressure differences can shield the line-driven wind from the central X-rays. Simulations such as those presented in \\cite{proga_stone_kallman} and  \\citet[][hereafter PK04]{pk04} also produce a shield, but in this case via a `failed wind'. This failed wind arises close to the X-ray-emitting region and quickly becomes over-ionized as it rises above the disk surface. As it falls back, it produces a dense shield that prevents the intense X-ray radiation from reaching the outer parts of the outflow in its shadow. These low-ionization regions are then able to interact with UV radiation and give rise to a strong line-driven wind. The same type of shielding structure has also been seen in other models of disk winds \\citep[e.g][]{ris_elv,nomura_13}.  PK04 used hydrodynamical simulations to investigate whether a line-driven outflow could be accelerated to high velocities purely by the radiation field produced by the accretion disk. They estimated the ionization state and temperature of the wind, taking into account only the central source of ionizing radiation, attenuated by electron scattering. Their main result was that the ionization state of the wind behind the failed wind region remained low enough to permit line driving, so that a fast, dense outflow was generated.  Absorption line profiles calculated from their wind model resembled those observed in BALQSOs (e.g. see Fig. 2 in in \\citealt{prog_kuro_10}). It is worth noting that PK04's wind, and especially its base, are virialized systems. This is because a line-driven wind accelerates slowly, so the rotational velocity is dominant over a relatively large distance while the wind base is Keplerian and very dense. This means that, if the broad line region (BLR) is associated with such a wind, black hole (BH) mass estimates based upon the assumption that the BLR is virialized remain correct \\citep{kashi_13}.  The PK04 wind model has already been subjected to two more detailed `post-processing' radiative transfer calculations. Both of these were concerned with the impact of the predicted outflow on the X-ray spectra of AGN/QSO. First, \\citet{schurch_09} performed a 1-D simulation which found that the wind produced observable X-ray spectral features. Second, \\citet[][hereafter SP10]{sim_proga_10} carried out a multi-dimensional radiative transfer simulation in which the ionization state and temperature structure of the wind were self-consistently computed. They confirmed that the wind was able to imprint a variety of characteristic features into the X-ray spectra of AGN. However, they also showed that scattering in the (failed) outflow is critically important in setting the ionization state of the wind in regions that would otherwise be shielded from the central engine.  PK04 used a simplified treatment of radiative transfer and ionization in their hydrodynamical simulations. These simplifications were essential to make the simulations computationally feasible. However, the work of SP10 implies that a careful treatment of radiative transfer may be required in order to obtain a reliable estimate of the ionization state of the wind. Since the ionization state, in turn, determines the efficiency of the line-driving mechanism itself, it is clearly important to check whether the wind model calculated by PK04 could actually maintain an ionization state that is consistent with efficient line driving and with the production of broad UV absorption lines.  Here, we therefore carry out a full, multi-dimensional radiative transfer calculation for the PK04 model that allows us to predict the ionization state, temperature structure and emergent UV spectra for this outflow. This extends the work done by SP10  to the longer wavelengths and lower ionization stages that are critical to allow effective line-driving and the formation of BALs. In particular, we account for the low-energy ($\\rm{< 0.1~keV}$) photons that affect the abundances of key ions, such as C~\\textsc{iv}, N~\\textsc{v} and O~\\textsc{vi}. We also calculate emergent spectra throughout the UV band, where the BAL features associated with the resonance transitions of these species are found. ",
        "conclusions": "We have presented results from comprehensive radiative transfer and photoionization calculations for the line-driven AGN disk wind predicted by the hydrodynamic simulations carried out in PK04. These simulations were computationally expensive and therefore treated the radiative transfer and ionization in a simplified manner. Here, we have focused on one snapshot from the hydrodynamic calculation and carried out a much more detailed Monte-Carlo simulation of the interaction between the AGN radiation field and the outflow.  Our main result is that the ionization state of the outflow is much higher than estimated by PK04, to the extent that line-driving would become inefficient. The over-ionized flow also no longer produces the broad UV absorption lines that are the key observational tracers of disk winds in AGN/QSOs. The main reason for this change in the predicted ionization state of the flow is that self-shielding becomes much less effective when radiative transfer effects are fully accounted for. More specifically, the failed wind region that protects the outflow from over-ionization in the simulations of PK04 can actually be ``circumnavigated'' by ionizing photons via scattering and reprocessing.  We conclude that hydrodynamic disk wind simulations need to take account of scattering and reprocessing in order to robustly assess the viability of line driving as an acceleration mechanism. This should ideally take the form of a self-consistent treatment, in which the radiative transfer and hydrodynamics are calculated simultaneously. However, the kind of `post-processing' approach we have used here is at least useful in validating purely hydrodynamic models. For outflows driven non-radiatively, it already provides a means to self-consistently predict the observational characteristics of the flow."
    },
    "1402/1402.4188_arXiv.txt": {
        "abstract": "Although the mechanisms responsible for heating the Sun's corona and accelerating the solar wind are still being actively investigated, it is largely accepted that photospheric motions provide the energy source and that the magnetic field must play a key role in the process. \\citet{2010ApJ...708L.116V} presented a model for heating and accelerating the solar wind based on the turbulent dissipation of Alfv\\'en waves.  We first use a time-dependent model of the solar wind to reproduce one of \\citeauthor{2010ApJ...708L.116V}'s solutions; then we  extend its application to the case when the energy equation includes thermal conduction and radiation losses, and the upper chromosphere is part of the computational domain. Using this model, we explore parameter space and describe the characteristics of a fast-solar-wind solution. We  discuss how  this formulation may be applied to a 3D MHD model  of the corona and solar wind \\citep{2009ApJ...690..902L}. ",
        "introduction": "The identification of the physical processes responsible for the heating of the solar corona and the acceleration of the solar wind still represents an unsolved problem in solar physics. However, there is a general consensus that photospheric motions provide the energy source and that the magnetic field must play a key role in the process. Since  it is clear that the measured speeds of fast streams require an extended heating deposition \\citep{1977ARA&A..15..363W,1980JGR....85.4665H,1988ApJ...325..442W}, previous one-dimensional (1D) models  generally relied on a parametric heating function exponentially decaying with height \\citep{1982ApJ...259..779H,1982ApJ...259..767H,1988ApJ...325..442W, 1995JGR...10021577H,1995GeoRL..22.1465R,1997ApJ...482..498H}. At the same time, several investigations were based on low-frequency broadband fluctuations on magnetohydrodynamic scales as the mechanism that heats and accelerates the solar wind \\citep{1968ApJ...153..371C,1971JGR....76.3534B,1986JGR....91.4111H, 1988JGR....93.9547H,1994AdSpR..14..123V,1999ApJ...523L..93M,2007ApJ...662..669V,2012ApJ...745...35Z}.   Connecting the macroscopic heating of the plasma and acceleration of the wind, as formulated in  coronal and inner heliospheric MHD models, with the underlying physical mechanisms is complicated due to the temporal and spatial dynamic ranges involved. In  the past decade, turbulent dissipation mechanisms have been progressively incorporated with various degrees of self-consistency into 1D models   of the  solar wind \\citep{2005ApJ...632L..49S,2005ApJS..156..265C,2007ApJS..171..520C,2010ApJ...710..676C, 2010ApJ...708L.116V,2011ApJ...743..197C}. There are also efforts to replace empirical heating functions in three-dimensional (3D) MHD models with some form of  turbulence dissipation mechanism. This is particularly challenging because such 3D models would have to resolve time-scales extending from a millisecond (dissipative time-scale in the solar corona) up to many days (large scale solar wind stream structure). \\citet{2010ApJ...725.1373V} introduced, beside the acceleration of the solar wind through Alfv\\'en waves, the  heating of the protons by Kolmogorov dissipation in open field-line regions. \\citet{2011ApJ...727...84U} developed a large-scale MHD heliospheric model with small-scale transport equations for the turbulence energy, normalized cross helicity, and correlation scale, applicable where the solar wind is already supersonic and superalfv\\'enic. The model of \\citet{2013AIPC.1539...30L} included  the effect of outwardly propagating Alfv\\'enic turbulence in the solar wind  and a phenomenological term to describe nonlinear interactions associated with wave reflection by density gradients in the chromosphere and corona. \\citet{2013ApJ...764...23S} and introduced Alfv\\'en wave turbulence, assuming that this turbulence and its nonlinear dissipation are the only momentum and energy source for heating the coronal plasma and driving the solar wind {\\citep[see][for additional details]{2013arXiv1311.4093V}. }  The present work illustrates the integration of  the turbulence dissipation heating and acceleration mechanism of \\citet{2010ApJ...708L.116V} in a time-dependent, 1D, hydrodynamic (HD) model of the solar wind, which includes also thermal conduction and radiation losses. In this early, explorative phase, it is expedient to conduct an investigation using 1D models so that we may later apply our gained experience to  3D models. The model of \\citet{2010ApJ...708L.116V} employs strong turbulence closure to treat nonlinear effects, and  does not rely on electron heat conduction for      radial energy transport, but rather computes the  internal energy associated with protons only. After presenting the characteristics of our {turbulence-driven HD} model, we show how it can match one of the solutions derived by \\citet{2010ApJ...708L.116V} without thermal conduction and radiation losses. Subsequently, we extend the application of our model to  include    transport mechanisms in the energy equation, which are necessary to reproduce plasma emission in agreement with observations \\citep{2009ApJ...690..902L}. Our exploration of parameter space yields solutions  compatible with the solar wind properties obtained from \\textit{in situ} measurements and observations. In the future we plan to introduce this formulation into the 3D MHD thermodynamic  model of the solar corona and solar wind of \\citet{2009ApJ...690..902L}.  This paper is organized as follows: the equations and the solution technique are described in Sec.~\\ref{sec-model}. In Sec.~\\ref{sec-results} we present  a solution in the configuration of \\citet{2010ApJ...708L.116V}, we conduct a parameter study that includes thermal conduction and radiative losses, and we describe the details of one of the solutions. We conclude with a discussion. ",
        "conclusions": "We have implemented a time-dependent model of the  solar wind with acceleration and heating through turbulence dissipation. Our model uses the self-consistent formulation of \\citet{2010ApJ...708L.116V} within a 1D HD code with thermal conduction and radiation losses. For the case of a simplified energy equation, such as that used by \\citet{2010ApJ...708L.116V}, we accurately reproduce a solar wind solution. When the model is extended by introducing thermal conduction and radiative losses, our model produces fast solar wind solutions whose characteristics are compatible with \\emph{in situ} measurements. The model of  \\citet{2010ApJ...708L.116V} includes a compressional heating in the lower corona; we have found that it is not necessary to include such phenomenological heating when a more realistic energy equation is used. Our model, in comparison with that of \\citet{2013AIPC.1539...30L}, is certainly more complicated, since it advances   the amplitude of the perturbations rather than the  energies as the latter does. Moreover, the model of \\citet{2013AIPC.1539...30L} neglects the contribution of the Reynolds stress, assuming that outwardly propagating wave is dominant. However, with the present method, we have found that the contribution of the Reynolds stress can be  sizable; in one particular case it lowered the wind speed by approximately $25\\%$.  Furthermore, if we consider the integration of turbulence dissipation heating and acceleration in 3D MHD models, there are several advantages in the present formulation in respect of that of  \\citet{2013AIPC.1539...30L}. First,  it is not much more computationally demanding to advance the amplitudes rather the energies. Second, this formulation in terms of $z_\\pm$ does not require us to calculate  for each mesh point the reflection coefficient along each field line passing through it. {Third, in our experience the inclusion of time-dependent transport equations for $z_\\pm$ does not increase the physical  convergence time to steady state.} Finally, as shown by \\citet{2011ApJ...727...84U}, the Reynolds stress term can be written for a case when only perturbations perpendicular to the large scale magnetic field are considered and whose actual directions do not play a significant role. Therefore we believe that the present formulation can be readily extended to the 3D model of \\cite{2009ApJ...690..902L}."
    },
    "1402/1402.1078_arXiv.txt": {
        "abstract": "We report the timing and spatial resolution from the Muon Telescope Detector (MTD) installed in the STAR experiment at RHIC. Cosmic ray muons traversing the STAR detector have an average transverse momentum of 6 GeV/$c$. Due to their very small multiple scattering, these cosmic muons provide an ideal tool to calibrate the detectors and measure their timing and spatial resolution. The values obtained were $\\sim$100 ps and $\\sim$1-2 cm, respectively. These values are comparable to those obtained from cosmic-ray bench tests and test beams. ",
        "introduction": "Data taken over the last decade have demonstrated that RHIC has created a hot, dense medium with partonic degrees of freedom called the Quark-Gluon Plasma (QGP). One of the physics goals for the next decade is to study the fundamental properties of the QGP such as the temperature, density profile, and color-screening length via electro-magnetic probes such as di-leptons~\\cite{starwhitepaper,rhicwhitepaper,dilepton,dileptonII, highptjpsi,satz_0512217,rhicIIQuarkonia,colorscreen,petreczky}. Muons can be measured more precisely because of their relatively reduced Bremsstrahlung radiation in the detector materials. Such an improved measurement is essential for separating the $\\Upsilon$ meson ground state (1S) from its excited states (2S+3S), each of which is predicted to melt at very different temperatures. The Muon Telescope Detector (MTD) in the Solenoidal Tracker at RHIC (STAR) will allow the measurement of the $\\Upsilon$ mesons and $J/\\psi$ mesons, over a broad transverse momentum range through di-muon decays to study color screening features, and $\\mu$-e correlations to distinguish heavy flavor correlations from initial lepton pair production~\\cite{starmtdproposal}. The MTD will thus provide direct information on the temperature and the characteristics of color screening in the QGP created in RHIC collisions.  The MTD is based on the Multi-gap Resistive Plate Chambers (MRPC) technology. A similar technology was used for the recently installed STAR TOF system~\\cite{startofproposal,startof}. Unlike the TOF MRPCs~\\cite{startofmrpcs}, however, the MTD MRPCs are much larger, and have long double-ended read-out strips. The MTD detectors are positioned behind the iron return bars of the STAR magnet, and cover 45\\% of the full azimuth within a pseudo-rapidity range of $|\\eta|$$<$0.5. The construction and installation of the MTD system was begun in 2011 and will be completed in 2014.  Prototype MTD detectors were built and studied from 2007 to 2011. These detectors were tested in the laboratory with cosmic rays, in test beams, and in STAR experiment during RHIC runs. The cosmic-ray and beam tests indicated a timing ($\\le$100 ps) and the spatial resolution ($\\sim$1 cm) that would be sufficient to achieve the physics goals~\\cite{MTDNIMA}. The operation of the prototype MTD detectors in STAR in 2007-2008 demonstrated that clean muon identification could be achieved for muon transverse momenta above a few GeV/$c$. The bench and test-beam results, as well as detailed simulations, thus indicated that the MTD would provide important physics information on quarkonia and primordial di-lepton measurements at RHIC~\\cite{MTDPerformanceAtSTAR}. However, during the tests of the MTD prototypes in STAR in 2007-2008, the timing resolution was 200-300 ps, which was much worse than that observed in the previous cosmic-ray and test-beam studies. This poorer resolution resulted from the particular digitization electronics~\\cite{startrigger} that were used at that time and the long cables between the MTD detectors and the digitizers.  In 2009, new prototype MTD detectors were installed in STAR. For these prototypes, the simple on-board front-end electronics used previously \\cite{mtdelectronics} to drive long cables to the digitizers were replaced with the same electronics as are used in the STAR TOF system~\\cite{startof,tofelectronics}. The TOF electronics are based on the CERN HPTDC~\\cite{HPTDC} chip. In 2010, a cosmic-ray trigger based on the information from the MTD prototypes was implemented in the STAR trigger system. High energy ($\\sim$6 GeV) cosmic rays provide an excellent means to study the timing and spatial resolution as the smearing from multiple scattering in the detector materials is a relatively small effect.  In this paper, the timing and spatial resolution of the MTD MRPCs read-out by STAR TOF electronics for cosmic muons reconstructed in the STAR experiment in 2010-2011 will be reported. The performance observed during the operation of 10\\% of the full MTD system during the 2012 RHIC run will also be presented. This paper is arranged as follows. Section~\\ref{mrpc} describes the MTD MRPCs, and Section~\\ref{expr} describes the experimental aspects. The details of the data analysis and the MRPC performance are reported in Section~\\ref{results}. The summary and conclusions are then presented in Section~\\ref{concl}. ",
        "conclusions": "The time and spatial resolution of the STAR MTD system were obtained using cosmic-ray muons traversing the STAR detector. The relatively high momentum muons can be cleanly triggered upon and selected in the data, and allow studies of the time and spatial resolutions with relatively small contributions from multiple scattering in the STAR detector materials. The MTD resolution values observed in data sets spanning several years were 100~ps for the timing resolution and 1-2 cm for the spatial resolution."
    },
    "1402/1402.1764_arXiv.txt": {
        "abstract": "\\noindent  One of the principal discoveries in modern cosmology is that standard model particles (including baryons, leptons and photons) together comprise only 5\\% of the mass-energy budget of the Universe\\cite{planck}. The remaining 95\\% consists of dark energy and dark matter (DM). Consequently our picture of the universe is known as $\\Lambda$CDM, with $\\Lambda$ denoting dark energy and CDM cold dark matter.  $\\Lambda$CDM is being challenged by its apparent inability to explain the low density of DM measured at the centre of cosmological systems, ranging from faint dwarf galaxies to massive clusters containing tens of galaxies the size of the Milky Way. But before making conclusions one should carefully include the effect of gas and stars, which were historically seen as merely a passive component during the assembly of galaxies. We now understand that these can in fact significantly alter the DM component, through a coupling based on rapid gravitational potential fluctuations. ",
        "introduction": "\\begin{figure*} \\includegraphics[width=0.49\\textwidth]{heic0604a.jpg} \\includegraphics[width=0.49\\textwidth]{martin_fig14.pdf} \\caption{The left panel (by J. Gallagher) shows a composite image of M82 taken by the Hubble Space Telescope. Purple colours correspond to narrow-band H$\\alpha$ emission, allowing us to see recombining hydrogen in outflowing gas.  The right panel, from Martin et al\\protect\\cite{martin12}, shows a compilation of measured absorption line blue-shifts for cool gas as a function of the galaxy's star formation rate. Even dwarf galaxies with star formation rates under $1\\,\\Msol\\,\\yr^{-1}$ are able to support winds exceeding $100\\,\\mathrm{km\\,s}^{-1}$. The outflow rate of these winds is typically several times the instantaneous star formation rate of the parent galaxy.} \\label{fig:outflows} \\end{figure*}  The viability of the $\\Lambda$CDM picture of structure formation was first evaluated using computer simulations (allowing, for instance, neutrinos to be ruled out as the dominant component of dark matter\\cite{frenk85}). Gas cooling and star formation within DM halos is now the standard paradigm for the origin of galaxies\\cite{whiterees78}. The behaviour of DM can be simulated on computers by chunking a portion of the universe into ``particles'' and evolving. Since the particles interact only through gravity, these simulations are called collisionless.  Early attempts used just $30\\,000$ particles to follow large regions of the Universe.  Consequently one particle had the mass of a large galaxy -- even so, such simulations were expensive, taking 70 CPU hours on state-of-the-art 3 MHz facilities. Such calculations would now take a few minutes on a cellphone.  The growth of computing power and parallel capabilities meant that, by the 1990s, simulations became sufficiently powerful to make detailed predictions of the internal structure of halos in different cosmological scenarios. These simulations highlighted the universal nature of DM halos formed through collisionless collapse. The spherically-averaged density of halos is `cusped' at the centre (scaling approximately as $\\rho \\propto r^{-1}$), rolling to a steeper slope at larger radius (reaching $\\rho \\propto r^{-3}$); such behaviour is known as ``NFW'' after the authors of a pivotal paper\\cite{navarro96}.  At the same time, simulations started highlighting a number of deficiencies in the CDM scenarios. The most evident was the overabundance, by more than an order of magnitude, of small satellites\\cite{moore99,klypin99} compared to the number observed orbiting the Milky Way\\cite{mateo98} at the time. Worse, the simulations significantly over-predicted the density of DM at the centre of galaxies\\cite{moore98}. Increasingly precise observations of the rotation curves of field galaxies have confirmed this discrepancy\\cite{deblok08} (see \\S3).  Collisionless DM simulations have since reached maturity, with modern simulations using several billion resolution elements for just one Milky Way sized halo\\cite{aquarius08,ghalo}.  However to make predictions which are testable against observations of the real Universe, baryon physics must be introduced. (Here we are adopting the astronomical convention of referring to baryons and leptons collectively as `baryons'.) Because baryons dissipate energy and so collapse to smaller scales than DM, they constitute a sizeable fraction of the mass in the central regions of all but the faintest galaxies \\cite{bell01}. Moreover observational constraints on galaxy formation ultimately come from photons, which can only be sourced by baryons.  Accordingly much effort has recently been devoted to implementing gas hydrodynamics and a description of star formation within simulations\\cite{gnedin09,enzo11,ramses,keres12}.  The energy released by young stellar populations and active galactic nuclei into the surrounding intergalactic medium is critical for regulating star formation\\cite{whiterees78}. Without this energy, most of the gas becomes cold and dense, rapidly collapsing to form stars, contradicting observations. Processes providing the energy to halt collapse are collectively named `feedback' and include supernova winds, radiation from young stars, and radiation and heat from black hole accretion \\cite{croton06,bower06,Stinson06,hopkins12}. Including these effects has led to strides forward in forming realistic disk galaxies, reproducing the efficiency of star formation as a function of galaxy mass, and linking gas accretion and mergers to galaxy morphology\\cite{robertson06,keres05,dekel09}.  However, until recently any direct effect of the baryonic component on the DM was limited to a minor `adiabatic' correction\\cite{blumenthal86} (see box A). In other words, star formation (SF) processes resulted in `passive' changes to the galaxy population -- modulating the star formation rate without significant changes to the underlying cosmic DM scaffolding.  This picture has recently been subverted. Spectroscopic observations reveal the ubiquity of massive galaxy outflows driven by feedback, carrying significant gas mass away from star forming galaxies throughout cosmic history\\cite{shapley03,weiner09,martin12} (see Section \\ref{sec:evid-galaxy-outfl}).  It has slowly been realised that these directly observed processes have a non-adiabatic impact on the associated dark matter halos. The effect is to relieve discrepancies between baseline CDM simulations and the real Universe (discussed in Section \\ref{sec:evidence-cusp-core}).  The emerging understanding of these processes constitute the central part of this review (Section \\ref{sec:one-process-rule}).  \\begin{figure*} \\vchb{\\includegraphics[width=0.49\\textwidth]{oh_2010_figure.pdf}} \\vchb{\\includegraphics[width=0.49\\textwidth]{weighted_hist_gnfw.pdf}} \\caption{The left panel is a compilation\\protect\\cite{oh11} of observed innermost dark matter density profile slopes ($\\alpha$ where $\\rho_{\\mathrm{DM}}(r) \\propto r^{\\alpha}$) for field dwarf galaxies, plotted at the innermost point where a robust determination has been achieved. Where the slope $\\alpha$ can be measured interior to around one kiloparsec, it is typically much shallower ($\\alpha>-1$) than the simulated ``NFW'' result. The right panel\\protect\\cite{newman13} shows the probability distribution function on the parameter $\\beta=-\\alpha$ for a selection of galaxy clusters ($M\\sim 10^{15}\\Msol$). While the constraints on individual clusters are quite broad, the combined constraints (thick line) again indicate a shallower-than-NFW slope. }\\label{fig:cores} \\end{figure*} ",
        "conclusions": "The $\\Lambda$CDM cosmology underlies a highly successful paradigm for explaining the formation of visible structure in the universe. Until recently, the key ingredients were passive processes which controlled the association of observable matter with the dark matter (for instance suppressing over-efficient star formation) while having little explicit effect on the underlying dark matter. There is, however, a new, rich literature of processes which violate this basic assumption and lead to fundamental modifications to the observable properties of galaxies. In the last few years these have come into sharp focus as increasingly sophisticated computer simulations have begun to follow the effects of star formation, and many relevant observational techniques have matured to the point that they can be regarded as robust.  Direct evidence of precisely which `baryonic processes' are in play and their relative importance in the real Universe at different scales should be our next priority. Because these baryonic processes simultaneously modify a number of observational diagnostics (outflows, dark matter cores, stellar morphology and star formation regulation), they weave into a coherent, testable framework.  It remains a possibility that tensions between observation and theory at the scale of faint dwarfs and clusters may point to exotic particle physics.  Ultimately we expect that a concerted effort from theorists and observers can achieve the goal of pointing to unique predictions of non-minimal DM models.  Of particular interest in the coming years will be ({\\it i}) improved understanding of the dark matter in dwarf spheroidals and faint field galaxies; if cores persist at the faintest end, it is a generic conclusion that baryonic physics cannot account for them\\cite{G12,penarrubia12}; ({\\it ii}) study of the stellar population ages and, separately, metallicity distributions of these objects to determine as far as possible whether the required bursty star formation histories are consistent propositions\\cite{mcquinn10}; ({\\it iii}) better predictions of the scalings of cores in massive galaxies and clusters for different scenarios; ({\\it iv}) observations that constrain the star formation histories of dwarfs\\cite{pacucci13} and the behaviour of gas at high redshift, especially through absorption line studies which are sensitive to internal kinematics and outflows\\cite{viel13}; ({\\it v}) renewed effort to understand how non-minimal dark matter scenarios (such as WDM or SIDM) interact with the revised, more complex baryonic physics of galaxy formation.  \\subsection{Acknowledgements}  We would like to thank Se-Heon Oh, Simon White, Max Pettini, Crystal Martin, Matt Walker, Jorge Pe{\\~n}arrubia, Alyson Brooks, Tommaso Treu, Richard Ellis, James Wadsley and Lisa Randall for helpful discussions and comments on an early draft."
    },
    "1402/1402.6691_arXiv.txt": {
        "abstract": "{Stellar differential rotation is important for understanding hydromagnetic stellar dynamos, instabilities, and transport processes in stellar interiors, as well as for a better treatment of tides in close binary and star-planet systems.}{We introduce a method of measuring a lower limit to the amplitude of surface differential rotation from high-precision, evenly sampled photometric time series,  such as those obtained by space-borne telescopes. It is designed to be applied to main-sequence late-type stars whose optical flux modulation is dominated by starspots.}{An autocorrelation of the time series was used to select stars that allow an accurate determination of  starspot rotation periods. A simple two-spot model was applied together with a Bayesian information criterion to preliminarily select intervals of the time series showing evidence of differential rotation with starspots of almost constant area. Finally, the significance of the differential rotation detection and a measurement of its amplitude and uncertainty  were obtained by an a posteriori Bayesian analysis based on a  Monte Carlo Markov Chain  approach. We applied our method to the Sun and eight other stars for which previous spot modelling had been performed to compare our results with previous ones.}{We find that autocorrelation is a  simple method for selecting stars with a coherent rotational signal that is a prerequisite for successfully measuring differential rotation through spot modelling.  For a proper  Monte Carlo Markov Chain analysis, it is necessary to take the strong correlations among different parameters that exist in spot modelling into account. For the planet-hosting star Kepler-30, we derive a lower limit to  the relative amplitude of the differential rotation of $\\Delta P / P = 0.0523 \\pm 0.0016$. We confirm that the Sun as a star in the optical passband is not suitable for measuring differential rotation owing to the rapid evolution of its photospheric active regions. In general, our method performs well in comparison to more sophisticated and time-consuming approaches.}{} ",
        "introduction": "\\label{intro}   The Sun and other stars do not rotate as rigid bodies owing to latitudinal and radial transport of angular momentum induced by anisotropic turbulent Reynolds stresses, meridional flows, and magnetic fields \\citep[e.g., ][]{Rudiger89}. Differential rotation (hereinafter DR) plays a fundamental role in  hydromagnetic dynamo \\citep{BrandenburgSubramanian05} and as a source of hydrodynamic and magnetohydrodynamic instabilities in stellar interiors \\citep{KnoblochSpruit82}. Moreover, it plays a role in tidal interaction in close binary systems \\citep{Scharlemann81,Scharlemann82} and is thus expected to affect the tidal interaction between a close-in planet and its host star \\citep{Mathisetal13}. The measurement of the mean rotation period of a main-sequence late-type star, from which its age is estimated by the method of gyrochronology, is also affected by the amplitude of latitudinal DR \\citep[see, e.g., ][]{EpsteinPinsonneault12}.  In the Sun, we can study DR in detail in the photosphere by measuring the rotation rate at different latitudes by  Doppler shifts of the plasma spectral lines as well as by using sunspots as tracers for the motion of the surrounding plasma. The interior DR is accessible by helioseismic techniques that reveal a time dependence in some of the layers, probably related to the feedback of the Lorentz force associated with hydromagnetic dynamo action \\citep[e.g., ][]{Howeetal00,Lanza07,Howe09}.  In distant stars, we have much more limited information because only spatially unresolved data can be acquired. Recently, asteroseismic techniques have provided the first hints on radial DR in red giants \\citep[e.g., ][]{Deheuvelsetal12,Mosseretal12,Goupiletal13}, while for main-sequence stars information on surface DR has been extracted through spectroscopic or photometric techniques. The advent of space-borne high-precision photometry with MOST \\citep[Microvariability and Oscillations of STars experiment, ][]{Rucinskietal03}, CoRoT \\citep[Convection, Rotation and Transit experiment, ][]{Auvergneetal09}, and Kepler \\citep{Boruckietal10} has made available large and homogeneous datasets of photometric measurements of late-type stars that represent a treasure trove to study stellar rotation and DR, in particular. Therefore, we introduce in the present work a technique to measure stellar DR from high-precision and evenly-sampled photometric time series of late-type main-sequence stars and discuss its advantages and limitations in the context of previously proposed approaches.  Main-sequence stars of the A and F spectral types are generally quite fast rotators and do not show brightness inhomogeneities in their photospheres { \\citep[see, however, ][]{Balona13}}, thus making  the effect of surface DR on the rotational broadening of spectral line profiles directly measurable by means of deconvolution techniques based on Fourier analysis \\citep[e.g., ][]{ReinersSchmitt03}. Applying this approach, \\citet{Reiners06} and \\citet{AmmlervonEiffReiners12} measured DR in a sample of A and F stars and found a remarkable increase in its amplitude with increasing effective temperature \\citep[see ][ for theoretical models that could explain such a dependence]{KukerRudiger05}. Stars of the spectral types G, K, and M generally show brightness inhomogeneities in their photospheres that are analogous to sunspots; i.e., they are associated with surface magnetic fields. Those having a sufficiently fast rotation ($ v \\sin i \\ga 15$~km~s$^{-1}$) can be mapped through the Doppler Imaging techniques \\citep[e.g., ][]{Donatietal97,DonatiCollierCameron97,Strassmeier09}  allowing  their surface DR to be measured  and its dependence on temperature and rotation rate to be studied. \\citet{Barnesetal05} find that the amplitude $\\Delta \\Omega$ of the DR has a weak dependence on the angular velocity of rotation $\\Omega$, i.e., $\\Delta \\Omega \\propto \\Omega^{\\alpha}$ with $\\alpha = 0.15 \\pm 0.10$, while a remarkably stronger correlation is found with stellar effective temperature, i.e.,  $\\Delta \\Omega$ increases with increasing effective temperature from M to G-type stars, thus extending the dependence found by \\citet{Reiners06} to lower temperatures. { In main-sequence stars, the detected differential rotation is solar-like; i.e.,  the equator rotates faster than the pole. Nevertheless, an anti-solar differential rotation has been suggested in some late-type giants \\citep[see, e.g., ][ and references therein]{Kovarietal07}}.  For late-type stars that are slowly rotating ($v \\sin i \\la 12-15$~km~s$^{-1}$), Doppler Imaging cannot be applied and information on surface DR can be extracted solely by photometric techniques. The chromospheric fluxes in the cores of the Ca~II~H\\&K lines have been monitored along several decades for a sample of late-type main-sequence stars in the framework of the Mt.~Wilson project to study rotation and stellar activity cycles. It has provided information on the dependence of  DR on stellar rotation rate \\citep{Donahueetal96} thanks to the quite long lifetime of chromospheric plages in comparison with the stellar rotation period that makes them good tracers for pointing out the differences in rotation rate at different latitudes \\citep{Donahueetal97a,Donahueetal97b}.  In the Sun, the activity belts where active regions form migrate with the phase of the activity cycle. This is interpreted as a migration of the latitude of maximum toroidal magnetic field close to the base of the convection zone \\citep[e.g., ][]{DikpatiCharbonneau99}. A similar migration is expected  in late-type stars that have a solar-like dynamo, producing a systematic variation in the period of the photometric modulation with the phase of the cycle. Such a variation has indeed been observed in the rotational modulation of the Sun-as-a-star chromospheric flux and provides an estimate of the amplitude of solar DR \\citep{DonahueKeil95}. A key parameter  is the length of the time interval used to determine the seasonal solar rotation period. It is calibrated by trying to match  two contrasting requirements: a) avoid remarkable variations of the large scale pattern of chromospheric inhomogeneities that would imply using as short an interval as possible; b) attain sufficient time resolution and low false-alarm probability in determining the period of the rotational modulation that would benefit from a time interval that is as long as possible. In the Sun, the optimal extension of the seasonal time interval is found to be $150-200$ days that makes a compromise between the two opposite requirements. This is allowed because chromospheric active regions and activity complexes are remarkably long-lived in comparison with photospheric spots having a mean lifetime of $50-80$ days vs. $\\sim 10-15$ days, respectively \\citep[cf. ][]{Donahueetal97a,Lanzaetal03}. As a matter of fact, a similar approach based on photospheric sunspots was not successful because of the random longitude appearance and short lifetime of individual sunspot groups \\citep{LaBonte82}.  The situation is different in the case of very active, young solar-like stars whose rotation period is shorter than the Sun's and whose photospheric starspots have lifetimes of several months. Therefore, the method was successful  in that case \\citep[see, e.g., ][]{MessinaGuinan02,MessinaGuinan03}. For the highly active and fast-rotating subgiant component stars in close binary systems such as RS Canum Venaticorum binaries, the persistence of active longitudes for decades allows us to measure a low-amplitude DR using photospheric starspots as tracers \\citep[$\\Delta \\Omega /\\Omega \\sim 10^{-3}$ in HR 1099, cf. ][]{Lanzaetal06,BerdyuginaHenry07}.  The majority of the stars observed by CoRoT and Kepler in the optical passband are not suitable to this approach because their photospheric active regions have lifetimes that are shorter than the typical timescale of DR shear, i.e., $1/\\Delta \\Omega$, where $\\Delta \\Omega$ is the amplitude of the DR. This limits the precision in the determination of the rotation period attainable with periodogram techniques, even in the case of a uniformly sampled time series \\citep{Lanzaetal93,Lanzaetal94}. On the other hand, if starspot intrinsic evolution is negligible, periodogram techniques coupled with a pre-whitening approach can be successful for estimating DR in solar-like stars by pinpointing the rotation frequencies of spots at different latitudes \\citep{ReinholdReiners13,Reinholdetal13}.  To make progress in the measurement of DR in late-type stars having starspots that evolve on a timescale comparable to $1/\\Delta \\Omega$ or possibly shorter, we investigate the potentiality of a simple starspot model to extract DR. Our approach applies a simple autocorrelation technique to estimate the coherence time of the light modulation that provides an estimate of the spot lifetime to be compared with the shear timescale. This allows us to select  promising candidates for spot modelling. For a given star, we perform a screening of the time intervals showing variations that  likely stem from the effect of DR rather than from intrinsic starspot  evolution. Finally, we apply a Monte Carlo Markov Chain (hereafter MCMC) method to estimate the most probable value of DR and its standard deviation following an  approach introduced by \\citet{Croll06}. We compare the proposed method with previous ones by analysing a sample of stars observed by MOST, CoRoT, and Kepler whose DR has been extracted with different spot modelling approaches. Moveover, we also consider the case of the Sun as a star to show the limitation of the method in the case of a slowly rotating star. ",
        "conclusions": "We have introduced a method of searching for differential rotation (DR) signals in high-precision  photometric time series of the late-type stars acquired by space-borne telescopes (see Sect.~\\ref{overview}). The even sampling of the time series  allowed us to apply an autocorrelation to measure the timescale of the intrinsic evolution of the spot pattern that limits the possibility of measuring DR. We selected a sample of eight stars for which previous determinations of  DR based on photometry were available along with some information on the inclination of the stellar rotation axis. We added the Sun as a star to this sample to study the rotation of solar analogues. We found that a significant DR can be detected  when  the relative height of the second maximum in the autocorrelation function is at least $0.6-0.7$.  In this case, we subdivided the time series into intervals the duration of which was shorter than the typical timescale of starspot evolution, although   long enough to reveal the drift of the spot longitudes produced by DR. The best interval was  then selected  by comparing the best fits obtained with a simple two-spot model with and without DR. For that interval, a fully Bayesian analysis was undertaken by means of an MCMC approach to assess the significance of DR and determine its most probable value and uncertainty.  We found that the amplitude of the DR derived with our method is generally consistent with previous results  when the different assumptions of the other approaches are considered. The advantage of our approach is the simplicity of the spot model that allows us to run MCMCs with tens or hundreds of million steps to study the a posteriori distribution of the parameters. This is particularly useful in spot modelling given the strong correlations among different parameters. We took advantage of the available information on the inclination of the stellar rotation axis to fix the a priori distribution of the inclination that is strongly correlated with the colatitudes and the unprojected areas of starspots in our model. Other correlations may become important when the signal-to-noise ratio of the photometry is low, spots are rapidly evolving ($\\tau_{\\rm s} \\la (2.5-3.0) P $), or the DR relative amplitude is small ($\\Delta P / P \\la 0.01$). They severely hamper the proper mixing and convergence of the MCMC procedure, although their effect can be controlled  as described in Sect.~\\ref{param_corr}, provided that a clear signal of DR is present in the light curve.  In the case of stars more active than the Sun, we find that the measured amplitude of DR depends on the specific time interval considered. In general, our approach only provides a lower limit to the amplitude of surface DR. We did not attempt to extract the amplitude and sign of the pole-equator shear as in other works \\citep[e.g., ][]{Froehlichetal12} because spot colatitude and inclination of the stellar spin axis are highly correlated in our simple spot model. On the other hand, an independent estimate of the inclination is generally not available, except in the case of close eclipsing binaries or transiting star-planet systems."
    },
    "1402/1402.4836_arXiv.txt": {
        "abstract": "We present updated analyses of pulse profiles and their arrival-times from PSR B1534+12, a 37.9-ms radio pulsar in orbit with another neutron star. A high-precision timing model is derived from twenty-two years of timing data, and accounts for all astrophysical processes that systematically affect pulse arrival-times. Five ``post-Keplerian\" parameters are measured that represent relativistic corrections to the standard Keplerian quantities of the pulsar's binary orbit. These relativistic parameters are then used to test general relativity by comparing the measurements with their predicted values. We conclude that relativity theory is confirmed to within 0.17\\% of its predictions. Furthermore, we derive the following astrophysical results from our timing analysis: a distance of $d_{\\textrm{GR}} = 1.051 \\pm 0.005$ kpc to the pulsar-binary system, by relating the ``excess\" orbital decay to Galactic parameters; evidence for pulse ``jitter\" in PSR B1534+12 due to short-term magnetospheric activity; and evolution in pulse-dispersion properties. As a secondary study, we also present several analyses on pulse-structure evolution and its connection to relativistic precession of the pulsar's spin axis. The precession-rate measurement yields a value of $\\Omega_1^{\\textrm{spin}}$ = $0.59^{+0.12}_{-0.08}$ $^{\\circ}$/year (68\\% confidence) that is consistent with expectations, and represents an additional test of relativistic gravity. ",
        "introduction": "Pulsars in relativistic binary systems have provided the most rigorous tests of gravitational theory in strong fields to date. High-precision timing of such an object produces a timing model that describes ``post-Keplerian\" (PK) effects that characterize relativistic corrections to the standard orbital elements \\citep{dd85,dd86}, as well as its nominal spin, astrometric, and environmental properties. Comparisons between measured and expected PK parameters produce tests of the gravitational theory in question. The ``Hulse-Taylor\" pulsar \\citep{ht75a} provided the first such positive case for general relativity, and still serves as an excellent laboratory for strong-field gravity \\citep{wnt10}. The recent discovery of a massive pulsar in a highly relativistic orbit with a white dwarf yields a testing ground for tensor-scalar extensions of gravitational theory \\citep{afw+13}. An extensive analysis of the ``double-pulsar\" system \\citep{bdp+03} constrains general relativity to within 0.05\\% of its predictions and remains the most stringent pulsar-timing test so far \\citep{ksm+06}.  PSR B1534+12 was discovered by \\citet{wol91a} to be in a highly inclined, 10.1-hour binary orbit with another neutron star. Follow-up timing studies on this pulsar \\citep{sac+98,sttw02} produced a timing solution that yielded measurements of five PK parameters: \\(\\dot{P}_b\\), the orbital decay of the binary system; \\(\\dot{\\omega}\\), the advance of periastron longitude; \\(\\gamma\\), the time-averaged gravitational-redshift and time-dilation parameter; \\(r\\textrm{ and }s\\), the ``range\" and ``shape\" of the Shapiro time delay. Simultaneous measurement of these parameters produced a self-consistent set of tests that complemented the Hulse-Taylor results by including tests based only on quasi-stationary, non-radiative PK parameters \\citep{twdw92}. Additional results were derived using the fitted timing model, including a precise estimate of the pulsar's distance using the measured excess of orbital decay due to relative motion \\citep{bb96}.    \\begin{deluxetable*}{ccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Timing Parameters for Each Backend and Frequency} \\tablewidth{0in} \\tablehead{ \\colhead{Parameter} & \\colhead{Mark III} & \\colhead{Mark IV} & \\colhead{Mark IV} & \\colhead{ASP} & \\colhead{ASP} & \\colhead{ASP} & \\colhead{ASP} & \\colhead{ASP} } \\startdata Frequency (MHz) \\dotfill & 1400 & 430 & 1400 & 424 & 428 & 432 & 436 & 1400 \\\\ Bandwidth (MHz) \\dotfill & 40 & 5 & 5 & 64 & 64 & 64 & 64 & 64 \\\\ Spectral Channels\\dotfill & 32 & 1\\tablenotemark{a} & 1\\tablenotemark{b} & 1 & 1 & 1 & 1 & 16\\tablenotemark{c} \\\\ Number of TOAs \\dotfill & 1185 & 3102 & 664 & 1204 & 1197 & 1190 & 1124 & 231 \\\\ Dedispersion type \\dotfill & Incoh. & Coh. & Coh. & Coh. & Coh. & Coh. & Coh. & Coh. \\\\ Integration time (s) \\dotfill & 300 & 190 & 190 & 180 & 180 & 180 & 180 & 180 \\\\ Date span (years) \\dotfill & 1990-94 & 1998-2005 & 1998-2005 & 2004-12 & 2004-12 & 2004-12 & 2004-12 & 2004-12 \\\\ RMS residual, \\(\\sigma_{rms}\\) & 5.31 & 4.21 & 6.73 & 4.48 & 4.34 & 4.65 & 4.93 & 8.27 \\enddata \\label{tab:backend} \\tablenotetext{a}{Four sub-bands centered at 430 MHz were taken when the Mark IV data were originally recorded, but were averaged together to build signal strength.} \\tablenotetext{b}{Two sub-bands centered at 1400 MHz were taken when the Mark IV data were originally recorded, but were also averaged together to build signal strength.} \\tablenotetext{c}{The number of actual channels recorded sometimes varied due to computational limitations, so this value represents a typical number of channels used.} \\end{deluxetable*}  The time-averaged pulse profile of PSR B1534+12 is undergoing a secular change in observed radiation pattern at a rate of 1\\% per year \\citep{arz95}. Such changes can be linked to spin-orbit coupling in a strong gravitational field, which results in a precession of the pulsar's spin axis  \\citep[``relativistic spin precession\";][]{ds16} and an evolving view of the two-dimensional beam structure \\citep{kra98}. \\citet*{sta04} (hereafter STA04) developed a general technique to characterize the overall profile shape at a given epoch and derive a precession rate by measuring and comparing spin-precession and orbital-aberration effects that produce the observed shape evolution. The results of this study yielded a direct measurement of the precession rate that was consistent with the rate predicted by general relativity, albeit with considerably limited precision. Furthermore, the geometry of PSR B1534+12 and its binary system was derived by combining these results with a rotating-vector-model  \\citep[RVM,][]{rc69a} analysis of the evolving polarization properties. PSR B1534+12 currently  remains the only pulsar for which special-relativistic orbital aberration is observed. {\\bfref The effects of relativistic spin precession on pulse structure have also been observed in PSR B1913+16 \\citep{wrt89}, the double-pulsar system \\citep{bkk+08}, and most dramatically in PSR J1141-6545 \\citep{mks+10}.}   In this work, we report on updated timing and profile-evolution analyses of PSR B1534+12, using data sets that collectively span 22 years since its discovery. Results from the analyses described below include improvements in tests of general relativity, an improved measurement of the pulsar's precession rate, and additional findings extracted from our time series. A full discussion of all current results is provided in Section \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Pulse Jitter in PSR B1534+12} \\label{sec:jitter}  The measurability of $\\ddot{\\nu}$ and $\\dddot{\\nu}$ in PSR B1534+12 strongly suggests a significant amount of timing noise across our data set. This polynomial whitening in our best-fit model removes most of the long-term timing noise, which is usually attributed to rotational instabilities and variations in magnetospheric torque. However, TOA residuals generally exhibit scatter on shorter, pulse-period timescales in excess of standard measurement uncertainties. This residual ``jitter\" is manifested from slight changes in the shape, amplitude and pulse phase of recorded profiles between successive pulses, and is likely due to variable activity within the pulsar magnetosphere. Recent studies of pulse jitter suggest that timing precision can be improved when averaging consecutive TOA residuals together \\citep{sc12}.  We believe that such pulse jitter is evident in our timing analysis of PSR B1534+12. Figure \\ref{fig:jitter} displays TOA residuals as a function of orbital phase from a global fit using unweighted TOA uncertainties recorded simultaneously on MJD 53545. The dark-blue data points were recorded with the Mark IV backend, while additional points represent the four channelized ASP data sets. While residual variations are visibly uncorrelated within the shown timescale, there is a visible correlation between the two overall data sets despite significant differences in backend specifications. We therefore associate this backend-correlated scatter as pulse jitter due to irregular activity within the pulsar's magnetosphere.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f4.pdf} \\caption{Pulse jitter in PSR B1534+12. The above figure shows global-fit residuals of TOAs recorded on MJD 53545 as a function of orbital phase. The dark-blue points were recorded with the Mark IV pulsar backend, and the remaining colors represent the four channels of ASP 430-MHz data recorded simultaneously. } \\label{fig:jitter} \\end{center} \\end{figure}  Several implications arise from the observed pulse jitter. First, pulse jitter will become a significant source of timing error in future timing studies of PSR B1534+12. The recent installation of the PUPPI signal processor\\footnote{\\url{http://www.naic.edu/~astro/guide/node11.html}} is expected to produce high-precision residuals with scatter that strongly reflects time-dependent inhomogeneities in the magnetosphere. Second, a long-term solution to jitter with PSR B1534+12 cannot involve averaging a large number of consecutive TOAs. The main objective of strong-gravity tests with pulsars is to monitor time-dependent changes to orbital elements and quasi-static PK parameters, which requires full coverage of the orbit over long periods of time. We therefore use this jitter as a means to justify the TOA-uncertainty compensation for the global timing solution described in Section \\ref{sec:obsred}. Lastly, further instrumental upgrades will only improve measurements made at 1400 MHz, where PSR B1534+12 is intrinsically weaker and signal-to-noise is currently limited. However, the overall timing solution is still expected to improve with additional observations over time.  \\begin{deluxetable*}{lcc} \\tabletypesize{\\scriptsize} \\tablecaption{Orbital Elements for PSR B1534+12} \\tablewidth{0in} \\tablehead{\\colhead{Parameter} & \\colhead{DD Model} & \\colhead{DDGR Model}} \\startdata Projected semimajor axis, \\(x\\) (s) \\dotfill & 3.7294636(6) & 3.72946417(13) \\\\ Eccentricity, \\(e\\) \\dotfill & 0.27367752(7) & 0.27367740(4) \\\\ Epoch of periastron, \\(T_0\\) (MJD) \\dotfill & 52076.827113263(11) & 52076.827113271(9) \\\\ Orbital Period, \\(P_b\\) (days) \\dotfill & 0.420737298879(2) & 0.420737298881(2) \\\\ Argument of periastron, \\(\\omega\\) (deg) \\dotfill & 283.306012(12) & 283.306029(10) \\\\ & & \\\\ Rate of periastron advance, \\(\\dot{\\omega}\\) (deg yr\\(^{-1}\\)) \\dotfill & 1.7557950(19) & 1.755795\\tablenotemark{a}  \\\\ Time-averaged gravitational redshift, \\(\\gamma\\) (ms) \\dotfill & 2.0708(5) & 2.0701\\tablenotemark{a} \\\\ Orbital decay, \\((\\dot{P}_b)^{\\textrm{obs}}\\) (\\(10^{-12}\\)) \\dotfill & -0.1366(3) & -0.19244\\tablenotemark{a} \\\\ Shape of Shapiro delay, \\(s = \\textrm{sin}i\\) \\dotfill & 0.9772(16) & 0.97496\\tablenotemark{a} \\\\ Range of Shapiro delay, \\(r = T_{\\odot}m_2\\) (\\(\\mu\\)s) \\dotfill & 6.6(2) & 6.627\\tablenotemark{a}  \\\\ & & \\\\ Companion mass, \\(m_2\\textrm{ }(M_{\\odot})\\) \\dotfill & 1.35(5) & 1.3455(2) \\\\ Pulsar mass, \\(m_1\\textrm{ }(M_{\\odot})\\) \\dotfill & $\\cdots$ & 1.3330(2)\\tablenotemark{a} \\\\ Total mass, \\(M = m_1+m_2 \\textrm{ }(M_{\\odot})\\) \\dotfill & $\\cdots$ &  2.678463(4) \\\\ Excess \\(\\dot{P}_b\\textrm{ }(10^{-12})\\) \\dotfill & $\\cdots$ & 0.0559(3) \\\\ \\enddata \\label{tab:binarypar} \\tablecomments{Values in parentheses denote the uncertainty in the preceding digit(s).} \\tablenotetext{a}{derived quantity} \\end{deluxetable*}    \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f5.pdf} \\caption{Phase structure function $D_{\\phi}$ as a function of time lag $\\tau$. The solid line is a best-fit model of Equation \\ref{eq:dmstructpred} for data with lags between 70 and 900 days.} \\label{fig:dmstruct} \\end{center} \\end{figure}  \\subsection{Long-term Variations in Dispersion Measure} \\label{sec:vardm}  Figure \\ref{fig:dm} illustrates an irregular evolution in DM over time for PSR B1534+12. Five large bins, each with a fitted gradient, are used in our timing model to fully describe the observed changes across different timespans. Recent studies have shown that several objects also exhibit nonlinear evolution in DM along different directions and distances \\citep[e.g.][]{kcs+13}. As with these studies, we believe that the dominant source of such evolution is the inhomogeneity of the interstellar medium, which is traced by the pulsar's signal as the line of sight sweeps through different regions due to a significant relative motion.  {\\bfref These long-term DM measurements are useful for a statistical analysis of turbulence within the interstellar medium \\citep*[e.g.][]{ktr94}, which usually assumes that the power spectrum of spatial variations in electron density is a power law within a range of length scales \\citep{ric90},  \\begin{equation} P(q) \\propto q^{-\\beta}, \\hspace{20pt} q_o < q < q_i \\label{eq:pq} \\end{equation}  \\noindent where $q = 2\\pi/l$ is a spatial frequency and $l$ is a scattering length. The frequency range in Equation \\ref{eq:pq} corresponds to a range between an ``inner\" ($l_i$) and ``outer\" ($l_o$) length scale where the power-law form is valid. The observed spatial fluctuations due to a relative transverse velocity $v$ are related to a time lag $\\tau$ by $l = v\\tau$. The power spectrum $P(q)$ can therefore be estimated by computing a pulse-phase structure function $D_{\\phi}(\\tau) = \\langle[\\phi(t+\\tau)-\\phi(t)]^2\\rangle$, where the angle brackets represent an ensemble average over observing epoch $t$. The pulse phase $\\phi$ is linearly related to DM, which therefore relates $D_{\\phi}(\\tau)$ to a DM structure function $D_{\\textrm{DM}}(\\tau) = \\langle[\\textrm{DM}(t+\\tau)- \\textrm{DM}(t)]^2\\rangle$,  \\begin{equation} D_{\\phi}(\\tau) = \\bigg(\\frac{2\\pi C}{f}\\bigg)^2D_{\\textrm{DM}}(\\tau) \\end{equation}  \\noindent where $C = 4.148\\times10^3 \\textrm{ MHz}^2\\textrm{ pc}^{-1}\\textrm{ cm}^3$ s, and $f$ is the observing frequency in MHz. Moreover, $D_{\\phi}(\\tau)$ is a power law in $\\tau$ within the inner length scales defined in Equation \\ref{eq:pq}, which finally requires that  \\begin{equation} D_{\\phi}(\\tau) = \\bigg(\\frac{\\tau}{\\tau_0}\\bigg)^{\\beta-2} \\label{eq:dmstructpred} \\end{equation}  \\noindent where $\\tau_0$ is a logarithmic intercept. Scintillation theory requires that $\\tau_0 = \\tau_\\textrm{d}$, where $\\tau_\\textrm{d}$ is the diffractive timescale, if the inner length-scale $l_i \\leq v\\tau_{\\textrm{d}}$.  We computed values of $D_{\\phi}(\\tau)$ at $f$ = 430 MHz using the Mark IV and ASP small-bin measurements of DM shown in Figure \\ref{fig:dm}. The Mark III DM point was measured using all Mark III TOAs collected over several years, which were generated with a different standard profile than the one used for the Mark IV and ASP data; we therefore chose to ignore this measurement in order to avoid incorporating bias in the structure function. Uncertainties in $D_{\\phi}(\\tau)$ were determined by propagating errors from our DM($t$) measurements. Our estimate of $D_{\\phi}(\\tau)$ is  shown in Figure \\ref{fig:dmstruct}, and illustrates a power-law evolution between time lags of roughly 70 and 900 days. We fitted Equation \\ref{eq:dmstructpred} to this segment of data, and found that  \\begin{align} \\beta &= 3.70 \\pm 0.04 \\nonumber \\\\ \\tau_0 &= 3.0 \\pm 0.8 \\textrm{ minutes} \\label{eq:n8} \\end{align}  \\noindent which is shown as a solid black in in Figure \\ref{fig:dmstruct}.  The measured spectral index $\\beta$ is consistent with the value for a ``Kolmogorov\" medium, $\\beta_{\\rm Kol}$ = 11/3. Furthermore, $\\beta$ and $\\tau_0$ in Equation \\ref{eq:n8} are consistent with the structure-function estimates reported by \\citet{sw12}. Our estimate of $\\tau_0$ is also consistent with the value of $\\tau_{\\textrm{d}}$ measured from the autocorrelation function of a dynamic spectrum of PSR B1534+12 \\citep{bplw02}.  At large timescales, the structure function departs from the fitted model at a lag $\\tau_o \\approx 900$ days, which suggests that  \\begin{equation} l_o \\approx 52 \\bigg(\\frac{v}{100\\textrm{ km/s}}\\bigg)\\textrm{ AU} \\label{eq:l_o} \\end{equation}  \\noindent \\citet{bplw02} derived an interstellar scintillation (ISS) velocity of 192 km/s. They noted in their study that ISS velocities of pulsars are dominated by the systemic transverse component, which means that $v \\approx 192$ km/s for PSR B1534+12, and $l_o \\sim 100$ AU $\\sim 10^{15}$ cm from Equation \\ref{eq:l_o}. This estimate is consistent with the upper limit of $l_o$ observed for several pulsars by \\citet{pw91}. By contrast, there is no evidence for a significant inner scale from our data set, since bins with mean values less than 70 days contain only one or two pairs of DM$(t)$ and were therefore ignored from the analysis. We did not apply any correction for the solar-wind contribution of our DM($t$) measurements, due to a covariance between the TEMPO solar-wind DM model and a fitted timing parameter that is discussed at the end of Section \\ref{sec:dist}.}  \\subsection{High-precision Distance to PSR B1534+12} \\label{sec:dist}  Relative acceleration between the observatory and pulsar systems in the Galactic potential causes significant Doppler-factor biases in PK parameters. \\citet{sac+98,sttw02} noted such behavior in earlier data sets of PSR B1534+12 with the observed orbital decay and applied a distance-dependent kinematic correction derived by \\citet{dt92} in order to include it as a consistent, radiative test of general relativity (see Section \\ref{sec:tests} below). However, the corrected value yielded a large uncertainty due to the imprecise distance to the pulsar derived from DM measurements using the \\citet{tc93} model of free electrons in the Galaxy. Previous studies of PSR B1534+12 therefore solved the inverse problem suggested by \\citet{bb96}, where general relativity is assumed to be correct; the distance is then derived using the measured ``excess\" orbital decay, a model of the Galactic acceleration \\citep{kg89}, and the expression for kinematic correction derived by \\citet{dt92}. Using this procedure, \\citet{sttw02} were able to derive a distance with a relative uncertainty of 4.9\\% when doubling their TEMPO uncertainties.  We used the same approach in this study to update the derived distance with a substantially longer timespan of Arecibo data. We also corrected the expression for the kinematic bias presented in \\citet{sac+98,sttw02} for missing factors of the cosine of the pulsar's galactic latitude; the correct equation is given by \\citet[Equation 5]{nt95}. Our derived distance to PSR B1534+12, using our timing results,  the corrected kinematic equation and updated Galactic parameters from \\citet{rmb+14}, is  {\\bfref \\begin{equation} d_{\\textrm{GR}} = 1.051 \\pm 0.005 \\textrm{ kpc} \\label{eq:distance} \\end{equation}   \\noindent where the value and its uncertainty (68\\% confidence level) were estimated using a Monte-Carlo method: all uncertain parameters were randomly sampled from a normal distribution with mean and standard deviation equal to their fitted values and uncertainties, respectively, which were then used to derive a value of $d_{\\textrm{GR}}$ by applying Newton-Raphson's method \\citep[e.g.][]{pftv86}. This process was repeated $10^5$ times, and resulted in a distribution of $d_{\\textrm{GR}}$ that is shown in Figure \\ref{fig:histdist}.} This new distance is consistent with the previous estimate of $1.02 \\pm 0.05$ kpc made by \\citet{sttw02}. The relative uncertainty of this result (0.48\\%) is slightly lower than that of the derived distance to PSR J0437-4715 estimated by \\citet{vbv+08}. We attribute this improvement in precision to the updated $(\\dot{P}_b)^{\\textrm{obs}}$ listed in Table \\ref{tab:binarypar}. The uncertainty in $d_{\\textrm{GR}}$ is dominated by uncertainties in Galactic acceleration and rotation parameters used to derive the estimate.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f6.pdf} \\caption{Distribution of $d_{\\textrm{GR}}$ obtained by using a Monte-Carlo method described in Section \\ref{sec:dist}.} \\label{fig:histdist} \\end{center} \\end{figure}  Despite its high level of precision, the derived distance presented in Equation \\ref{eq:distance} is a model-dependent quantity. An ideal measure of distance can be obtained from the geometric, model-independent parallax through low-frequency interferometry. The recent inclusion of PSR B1534+12 into an extension of the PSR$\\pi$ interferometry program\\footnote{\\url{https://safe.nrao.edu/vlba/psrpi}} will likely provide such an estimate in the next 2-3 years. Another independent distance measure can be derived from a timing parallax estimated in the global-fit timing solution. However, our measured timing parallax was found to be significantly covariant with an input DM parameter associated with free electrons from the solar wind. Such a covariance is unexpected since the solar-wind component of DM is strongest for pulsars close to the ecliptic plane, while PSR B1534+12 is $\\sim 30^{\\circ}$ above the plane. Since the expected solar contribution is much smaller than the scatter of the 80-day DM bins in Figure \\ref{fig:dm}, we chose to set the solar DM component to zero for the global timing fit while acknowledging that the timing parallax is unreliable as a fitted parameter.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f7.pdf} \\caption{\\emph{Top}: MCMC posterior distribution of $\\Omega_1^{\\textrm{spin}}$ obtained from the profile-shape analysis of Mark IV and ASP data discussed in Section \\ref{sec:precres}. \\emph{Bottom}: Markov chain for $\\Omega_1^{\\textrm{spin}}$ determined from the MCMC algorithm.} \\label{fig:mcmcpos} \\end{center} \\end{figure}  \\subsection{Precession and Geometry} \\label{sec:precres}  Results from the MCMC fit on several data sets can be found in Table \\ref{tab:profdata}, and the posterior distribution for $\\Omega_1^{\\textrm{spin}}$ derived from our Mark IV and ASP data sets is shown in Figure \\ref{fig:mcmcpos}. {\\bfref We generated $3\\times10^5$ samples for each application of the algorithm, after burning the first 5000 samples in order to remove non-convergent iterations.} We provided the original results obtained by STA04, as well as a reproduced set of results from the STA04 data set using the MCMC algorithm, for comparison with our extended Mark IV and ASP profiles. We assumed that values of $F'$ must be negative while using the MCMC algorithm, since the simultaneous-linear-fit technique used and described by STA04, which avoids any consideration of $F'$, estimates that $\\cos\\eta < 0$. These results agree well with predictions from general relativity, where $\\Omega_1^{\\textrm{spin}} = 0.51^{\\circ}$/yr using the derived masses in Table \\ref{tab:binarypar}, and previous measurements made by STA04. General improvements in precision come from the new fitting procedure, which permitted direct sampling of the precession rate and other free parameters, as well as the addition of the ASP 2005 campaign and several strong bi-monthly observations.  \\begin{deluxetable}{lcccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Profile-evolution MCMC Parameters} \\tablehead{\\colhead{Parameter} & \\colhead{STA04\\tablenotemark{a}} & \\colhead{STA04} & \\colhead{Mark IV} & \\colhead{All}} \\startdata $\\Omega_1^{\\textrm{spin}}$ $(^{\\circ}/\\textrm{yr})$ \\dotfill & $0.44^{+0.48}_{-0.16}$ & $0.51^{+0.10}_{-0.08}$ & $0.48^{+0.09}_{-0.07}$ & $0.59^{+0.12}_{-0.08}$ \\\\ [5pt] $\\eta$ $(^{\\circ})$ \\dotfill & $\\pm103^{+10}_{-10}$ & $\\pm99^{+2}_{-2}$ & $\\pm118^{+10}_{-15}$ & $\\pm139^{+16}_{-25}$ \\\\ [5pt] $F'$ \\dotfill & n/a & $-5.9^{+0.9}_{-1.0}$ & $-2.2^{+0.6}_{-0.7}$ & $-1.3^{+0.3}_{-0.5}$ \\\\ [5pt] $\\epsilon$ ($10^{-3}$) \\dotfill & $-1.5^{+0.3}_{-0.3}$ & $-1.90^{+0.08}_{-0.09}$  & $-1.21^{+0.08}_{-0.08}$ & $6.67^{+0.07}_{-0.07}$ \\\\ [5pt] \\enddata \\label{tab:profdata} \\tablenotetext{a}{Original, non-MCMC results from STA04.} \\tablecomments{Uncertainties reflect 68\\% confidence intervals of posterior distributions.} \\end{deluxetable}  The RVM analysis yielded values of $\\alpha$ and $\\beta$ at different times using the Mark III (reticon), Mark IV and ASP campaign profiles. The values of $\\beta$ measured for each campaign are shown in Figure \\ref{fig:betavtime}. Measurements of $\\alpha = 103.5(3)^{\\circ}$ are consistent with no evolution in time, while the values of $\\beta$ are found to change significantly, where $d\\beta/dt$ = -0.23 $\\pm$ 0.02 $^{\\circ}$/yr. This is consistent with the STA04 result of -0.21 $\\pm$ 0.03 $^{\\circ}$/yr. The assumption that general relativity is correct requires that $d\\beta/dt = \\Omega_1^{\\textrm{spin}} \\sin i\\cos\\eta$, and therefore yields $\\eta = \\pm 117 \\pm 3^{\\circ}$ (68\\% confidence), which agrees with the value determined from the MCMC analysis described above. With these values, the misalignment angle $\\delta$ between the spin and orbital angular momentum axes can be derived through  spherical trigonometry by $\\cos\\delta = -\\sin i\\sin\\lambda\\sin\\eta+\\cos\\lambda\\cos i$. The sign ambiguity in $\\eta$ and $i$, as well as the requirement that $\\cos i\\tan\\eta > 0$ pointed out by STA04, gives an expected value of $\\delta = 27.0 \\pm 3.0^{\\circ}$ or $\\delta = 153.0 \\pm 3.0^{\\circ}$. Physical arguments based on alignment of angular momenta prior to the second supernova suggest that the smaller angle is correct \\citep{bai88}, and therefore requires that $\\eta = -117 \\pm 3^{\\circ}$ and $i = 77.7 \\pm 0.9^{\\circ}$.  The consistency between the MCMC and RVM analyses serves as an improved, independent check of precession within this relativistic binary system. These results also confirm the geometric picture of this pulsar-binary system derived in STA04.  \\citet{kws03} pointed out that relativistic spin precession can eventually imprint a timing signature as a second-derivative in spin frequency. However, they predicted that it would take an additional 25 years of timing observations in order to measure the predicted signature with reasonable accuracy. Moreover, our best-fit timing model of PSR B1534+12 data indicates that the measured $\\ddot{\\nu}$ and $\\dddot{\\nu}$ are dominated by timing noise.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f8.pdf} \\caption{Impact angle $\\beta$ between the magnetic axis and line of sight as a function of time. The black line is a best-fit slope of -0.23 $\\pm$ 0.02 $^{\\circ}$/yr.} \\label{fig:betavtime} \\end{center} \\end{figure}  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.43]{./f9.pdf} \\caption{Mass-mass plot for PSR B1534+12. Each set of black curves represents the 68\\%-confidence region delimited by the labeled PK (timing) parameter. The dot-dashed red curves represent the 68\\% confidence region determined from the spin-precession rate. } \\label{fig:m1m2} \\end{center} \\end{figure}  \\subsection{Tests of General Relativity} \\label{sec:tests}  In general relativity, each PK parameter is expressed as a function of at least one of the two binary-component masses; one can therefore define a ``mass-mass\" space where each PK parameter corresponds to a curve in this plane.  \\citet{sac+98} presented the first mass-mass plot that incorporated up to 5 PK curves,  as well as the first ``non-mixed\" test of quasi-static parameters using the $\\dot{\\omega}-\\gamma-s$ combination.  Figure \\ref{fig:m1m2} presents our evaluation of PK parameters and tests of general relativity using PSR B1534+12. Each curve corresponds to a PK value measured either with the best-fit DD timing model or profile-shape model, while the filled circle represents the best-fit DDGR solution of the pulsar and companion masses of $m_1 = 1.3330\\textrm{ }M_{\\odot}$, $m_2 = 1.3455\\textrm{ }M_{\\odot}$, respectively. The $\\dot{P}_b$ curve was corrected for kinematic bias assuming a pulsar distance of $d = 0.7 \\pm 0.2$ kpc, which was estimated using the electron number-density model developed by \\citet{tc93}, as discussed in Section \\ref{sec:dist}. The discrepancy between this distance and the expected distance (Equation \\ref{eq:distance}) prevents the curve from intersecting the other PK curves, while its large uncertainty dominates the corrected orbital-decay error estimate. The $\\Omega_1^{\\textrm{spin}}$ curves intersect the best-fit DDGR point well at the 68\\% confidence level.  The Shapiro $r$ parameter remains a slightly weaker constraint than the PK timing parameters, but its relative uncertainty has improved by nearly a factor of two since the last measurement by \\citet{sttw02}. The $\\dot{\\omega}-\\gamma-s$ combination is the strongest test from this pulsar and confirms general relativity to within 0.17\\% of its predictions. This test quality is slightly larger than the 0.05\\% test from the double-pulsar system, which uses the mass ratio as determined by the projected semi-major axes of both pulsars \\citep{ksm+06}.  The high-precision DDGR masses of PSR B1534+12 and its companion are consistent with previous estimates. As noted in \\citet{sttw02}, the significant difference between $m_1$ and $m_2$ presents a conundrum where the spun-up pulsar is actually less massive than its companion. This suggests that a period of ``mass inversion\" -- mass transfer from the pulsar's progenitor to its companion -- took place during the system's evolution, though a more thorough understanding of mass-transfer processes and its effects on stellar structure is needed. The effectiveness of mass estimates from Shapiro-delay measurements will continue to provide a better view of the general pulsar-companion mass population \\citep{kkdt13}.\\\\   The Arecibo Observatory is operated by SRI International under a cooperative agreement with the National Science Foundation (AST-1100968), and in alliance with Ana G. M\\'{e}ndez-Universidad Metropolitana, and the Universities Space Research Association. Pulsar research at UBC is supported by an NSERC Discovery Grant. We thank Z. Arzoumanian, F. Camilo, A. Lyne, D. Nice, J. H. Taylor, and A. Wolszczan for their earlier contributions to this project. We also thank P. Freire, I. Hoffman, A. Lommen, D. Lorimer, D. Nice, E. Splaver, and K. Xilouris for previous assistance with Mark-IV observations. We also thank R. Ferdman, M. Gonzalez, other NANOGrav observers for help with some of the ASP observations reported here. We thank W. Zhu for providing a TEMPO MCMC program. We are grateful for useful comments provided by an anonymous referee."
    },
    "1402/1402.6002_arXiv.txt": {
        "abstract": "The Canadian Automated Meteor Observatory (CAMO) detects occasional meteors with two maxima in the image intensified CCD based light curves. We report early results from an analysis of 21 of these events.  Most of these events show qualitatively similar light curves, with a rounded first luminous peak, followed by an almost linear sharp rise in the second peak, and a relatively rapid curved decay of the second peak. While a number of mechanisms could explain two maxima in the light curves, numerical modelling shows that most of these events can be matched by a simple dustball model in which some grains have been released well before intensive ablation begins, followed by a later release of core grains at a single time.  Best fits to observations are obtained with the core grains being larger than the pre-released outer grains, with the core grains typically $10^{-6}$~kg while the early release grains are of the order of $10^{-9}$~kg. ",
        "introduction": "Light curves provide one of the best indicators of the mode of meteor  ablation and fragmentation, and implied meteoroid structure.  The intensity of the light produced is indicative of the almost instantaneous mass loss rate.  In this paper we consider implications for meteoroid structure and ablation from dual peak meteor light curves, in particular looking at whether these observations support a dustball model.  \\cite{Jacchia1955} suggested that photographic observations of shortened trails, flares and meteor wake supported a dustball structure for at least some meteors.  However the fragile dustball structure proposed was inconsistent with bright meteors that survived high pressures to low observed heights.  \\cite{HawkesJones1975} developed a two component dustball model that would fit both faint and brighter meteor light curves. The key idea of this model was that meteoroids have two components, a grain component with a silicate metallic composition that is responsible for the light production, and a glue component (possibly organic in nature) which has a lower boiling point and does not produce significant luminosity.  With this model meteoroids are not necessarily fragile.  Under this model some grains will be released early (as the glue reaches its boiling point) while other grains will be separated during intensive light production by the meteor.  \\cite{FisherEt2000} reviewed the observational evidence in favour of this dual component dustball model (e.g. shorter more symmetric light curves, relative independence of heights of meteors past a certain mass cutoff, good prediction of ablation profiles across a wide mass range). ",
        "conclusions": "Perhaps the most obvious question is why some meteors, but not all, show a dual peak structure.  A related question is why these events have not been more frequently reported in the past.  One possibility is that the relatively high response in the near infrared of the CAMO  detection systems may play a factor.  One of the most striking aspects of the study was the similarity of  different light curves. In almost all cases the first peak is more rounded and symmetric, while the second peak has a sharp almost linear (in a plot using the logarithmic astronomical magnitude as the vertical axis), followed by a fairly sharp rounded decline. This suggests to us that a similar mechanism is producing most of the double light curve events.  While other mechanisms can probably also match the results, the simple dustball model used here seems adequate for the events studied. \\cite{MalhotraMathews2010} have conducted a statistical study of smaller meteoroids observed with large aperture radar, and find that only about one-quarter are consistent with simple single body ablation. They find that 48\\% seem to show fragmentation, while 20\\% imply differential ablation is important.  If a more detailed future analysis continues to suggest that the core grains are larger than the early release grains, that is an important and somewhat surprising result. This work suggests that core grains  are typically~$10^{-6}$~kg while  early release grains are of the order of ~$10^{-9}$~kg.  A number of previous studies have sought to establish the size of meteor constituent grains. \\cite{Simonenko1968} found from an analysis of rapid onset flares on bright meteors a mean grain mass of about  $2\\times10^{-9}$~kg.  \\cite{BorovickaEt2007} have modelled six electro-optical Draconid meteors, using an erosion model in which grains are released over generally the first half of the light curve, although for the brightest meteor a number of grains were resistant to release and separated later. They generally found agreement using grains from $10^{-11}$ to $10^{-9}$ kg. They find a model with total pre-release prior to intensive ablation inconsistent with the meteors in this small sample. Of course it should be kept in mind that these were Draconid meteors, known to have a porous low density structure (\\cite{BorovickaEt2007}).  \\cite{CampbellKoschny2004} have modelled Leonid light curves using a thermal erosion two component dustball model. The  find that two of the meteors are well matched with Gaussian mass distribution of grains, while the third requires a power law distribution.  Grain sizes generally range from $1\\times10^{-11}$ to $4\\times10^{-7}$~kg, in general agreement with this study.  \\cite{BeechMurray2003} used a power-law mass distribution to match Leonid light curves, generally finding the need for grains in the mass range from $10^{-10}$ to $10^{-7}$~kg. In an analysis of four Leonid fireball bursts \\cite{HawkesEt2002} found that relatively large grains $10^{-5}$ to $10^{-4}$~kg were needed to match the observations.  It is also interesting to compare the grain size distribution with that obtained by Stardust during encounter with the coma of Comet 81P/Wild~2. \\cite{Green2007} show that most grains are smaller than those reported here, although the total range is about $10^{-15}$ to $2\\times10^{-5}$~kg.  It is possible that the outer grains released prior to ablation are produced in a similar manner to the less spatially and temporally constrained clusters of meteoroids occasionally observed (\\cite{WatanabeEt2003, PiersHawkes1993}). While \\cite{WatanabeEt2003} discusses possible production mechanisms, there is not yet clarity as to how these clusters occur.  \\cite{BroschEt2004} point out the many parameters of meteor light curve analysis, and how they can help to constrain meteoroid structure and ablation. \\cite{StokanEt2013} have studied optical trail widths using the CAMO system and this information will help constrain the possible meteoroid models.  We plan to in a later paper incorporate the tracking data for a small number of events in a more in depth investigation of dual light curve events which may help to further constrain possible models."
    },
    "1402/1402.3627_arXiv.txt": {
        "abstract": "First evidence of high-energy astrophysical neutrino observation with the IceCube detector from May 2010 to May 2012 is presented. Selecting for high-energy neutrino events with vertices well contained in the detector volume, the analysis has sensitivity starting at approximately 50 TeV, and the highest energy events are a few PeV. A significant excess in flux is observed above expected backgrounds from atmospheric muons and neutrinos. The sample of 28 events includes the highest energy neutrinos ever observed, and has properties consistent in flavor, arrival direction, and energy with generic expectations for neutrinos of extraterrestrial origin. Using the new found astrophysical neutrinos, spatial searches of clusters are performed to look for exact sources, marking the birth of neutirno astronomy. ",
        "introduction": "Observation of high-energy neutrinos are thought to provide insight into the origins and acceleration mechanisms of high-energy cosmic rays. Cosmic-ray protons and nuclei produce neutrinos in interactions with gas in their source environment and interstellar dust through decay of charged pions and kaons. These neutrinos have energies proportional to the cosmic rays that produced them and point back to their sources since they are neither affected by magnetic fields nor absorbed by matter opaque to radiation. Neutrinos can also provide insight into the creation of gamma-rays at astronomical sources. Gamma-rays can be produced by mesons from cosmic ray interactions in a similar process to neutrinos. This is known as the hadronic process, but they can also be produced by leptonic processes such as inverse compton scattering of electrons in high-energy source environments. Neutrino observation of gamma-ray sources will confirm their hadronic origin. Large-volume Cherenkov detectors like IceCube \\cite{daqpaper} can detect these neutrinos through production of secondary leptons and hadronic showers when they interact with the detector material. This analysis has an event selection with an energy threshold at about 50 TeV and is sensitive to energies up to 10 PeV and beyond, depending on the flux being measured. By selecting for events well contained in the detector volume, it is sensitive to all neutrino flavors from all directions. Hints of astrophysical sources are seeked by characterizing the astrophysical flux measured and looking for spatial clustering among these events.  The analysis uses a data-taking period that started in May 2010 using 79 strings and continuing with the completed detector (86 strings) from May 2011 to May 2012 for a total livetime of 662 days. ",
        "conclusions": "An analysis of two years of IceCube data from 2010 to 2012 has observed a flux incompatible with expectations from terrestrial processes. It contains a mixture of neutrino flavors compatible with a flux proportional to $E^{-2}$, a spectrum expected for neutrinos associated with primary cosmic ray acceleration. The sample is thus consistent with generic expectations for a neutrino population with origins outside the solar system. We did not observe significant spatial clustering of the events, although this study is currently limited by low statistics and poor angular resolution for the majority of the observed events. Future observations with IceCube will provide improved measurements of the energy spectrum and origins of this flux, providing insight into the underlying processes responsible for these events."
    },
    "1402/1402.5976_arXiv.txt": {
        "abstract": "We present spectral energy distributions (SEDs) of 69 QSOs at $z>5$, covering a rest frame wavelength range of 0.1\\,$\\mu$m to $\\sim$80\\,$\\mu$m, and centered on new {\\it Spitzer} and {\\it Herschel} observations. The detection rate of the QSOs with {\\it Spitzer} is very high (97\\% at $\\lambda_{\\rm rest}\\,\\lesssim\\,4\\,\\mu$m), but drops towards the {\\it Herschel} bands with 30\\%  detected in PACS (rest frame mid-infrared) and 15\\% additionally in the SPIRE (rest frame far-infrared; FIR). We perform multi-component SED fits for {\\it Herschel}-detected objects and confirm that to match the observed SEDs, a clumpy torus model needs to be complemented by a hot ($\\sim$1300K) component and, in cases with prominent FIR emission, also by a cold ($\\sim$50K) component. In the FIR detected cases the luminosity of the cold component is on the order of $10^{13}$ $L_{\\odot}$ which is likely heated by star formation. From the SED fits we also determine that the AGN dust-to-accretion disk luminosity ratio declines with UV/optical luminosity. Emission from hot ($\\sim$1300K) dust is common in our sample, showing that nuclear dust is ubiquitous in luminous QSOs out to redshift 6.  However, about 15\\% of the objects appear under-luminous in the near infrared compared to their optical emission and seem to be deficient in (but not devoid of) hot dust. Within our full sample, the QSOs detected with {\\it Herschel} are found at the high luminosity end in $L_{\\rm UV/opt}$ and $L_{\\rm NIR}$ and show low equivalent widths (EWs) in H$\\alpha$ and in Ly$\\alpha$. In the distribution of H$\\alpha$ EWs, as determined from the {\\it Spitzer} photometry, the high-redshift QSOs show little difference to low redshift AGN. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.2290_arXiv.txt": {
        "abstract": "We explore the use of different radio galaxy populations as tracers of different mass halos and therefore, with different bias properties, to constrain primordial non-Gaussianity of the local type. We perform a Fisher matrix analysis based on the predicted auto and cross angular power spectra of these populations, using simulated redshift distributions as a function of detection flux and the evolution of the bias for the different galaxy types (Star forming galaxies, Starburst galaxies, Radio-Quiet Quasars, FRI and FRII AGN galaxies). We show that such a multi-tracer analysis greatly improves the information on non-Gaussianity by drastically reducing the cosmic variance contribution to the overall error budget. By using this method applied to future surveys, we predict a constraint of $\\sigma_{f_{nl}}=3.6$ on the local non-Gaussian parameter for a galaxy detection flux limit of 10$ \\mu$Jy and $\\sigma_{f_{nl}}=2.2$ for 1 $\\mu$Jy. We show that this significantly improves on the constraints obtained when using the whole undifferentiated populations ($\\sigma_{f_{nl}}=48$ for 10 $\\mu$Jy and $\\sigma_{f_{nl}}=12$ for 1 $\\mu$Jy). We conclude that continuum radio surveys alone have the potential to constrain primordial non-Gaussianity to an accuracy at least a factor of two better than the present constraints obtained with Planck data on the CMB bispectrum, opening a window to obtain $\\sigma_{f_{nl}}\\sim1$ with the Square Kilometer Array. ",
        "introduction": "In the current standard cosmological model, the large scale structures observed in the Universe originated from small fluctuations present in the matter density field arising after the inflationary phase shortly after the Big Bang. Although slow-roll inflation models predict this random field to be essentially Gaussian \\citep{Maldacena, Acquaviva}, other evolutionary models after inflation predict a non-vanishing non-Gaussian component in the primordial matter density field \\citep{Verde, Liguori}. The detection of non-Gaussianity could open a new window in the knowledge of early Universe physics.  The most widely used method so far to constrain primordial non-Gaussianity is to measure the bi-spectrum in the Cosmic Microwave Background (CMB) temperature anisotropy maps. Recently, this method was applied to Planck data \\citep{Planck2} to provide the best measurement of local non-Gaussianity up to date ($f_{nl}=2.7 \\pm 5.8$)\\footnote{Note that in our definition of $f_{nl}$, with the growth factor normalized to unity today, this value should be multiplied by a factor of $\\approx 1.3$ \\citep[see e.g.][]{Dalal,Afshordi}.}. A complementary way to access non-Gaussianity is to measure its impact on Large Scale Structure (LSS) at lower redshifts \\citep{Dalal,b2, carbone} which affects the bias of dark matter tracers. Large galaxy surveys have been the probe of choice to constrain the clustering properties of dark matter on large scales and measure this non-Gaussian effect \\citep{Xia2010, Xia2011, Bernardis}. Given the fact that this non-Gaussian signal is specially sensitive on large scales, other type of surveys, based on intensity mapping techniques, have been suggested to go after this effect \\citep{Joudaki, Camera2013}.  Although observing large scales has the advantage of probing the regime where the non-Gaussian effect on the bias of the dark matter tracers is stronger, it has the problem that it is nonetheless limited by cosmic variance, e.g., the lack of enough independent measurements for the scales we are trying to probe, given the limited size of the volume that is observed. \\citet{Seljak} proposed a way to get around this cosmic variance limitation by using different biased tracers of the underlying dark matter distribution. With at least two tracers with different bias, we can make a measurement of the ratio of these two biases that is only limited by shot noise and hence beats cosmic variance. This is specially sensitive when the bias of one object is much larger than the other. In order to understand this, let us assume that we measure two density maps at a given redshift, one for a biased tracer and another for the dark matter itself. Due to cosmic variance, there will be several power spectra (e.g. several cosmologies) that are consistent with the dark matter map. However, the ratio of the two maps should give a direct measurement of the bias, with an uncertainty just given by the shot noise of the tracer.  In this paper we propose to use the mass of the haloes hosting the dark matter tracers as an extra source of information in order to constrain the bias. This will have 3 advantages: 1) it will allow us to select objects with large bias factors (without any mixture with low mass objects) making them more sensitive to the non-Gaussian effect; 2) it will allow us to compare directly the bias of different tracers in order to cancel cosmic variance and 3) it will allow us to have a more physical description of the bias parameters used in the analysis. Although this idea can be applied to different surveys, it can be particularly relevant to radio galaxy surveys, given that they usually lack redshift information but on the other hand should provide a more direct relation between bias and mass. In here we concentrate on surveys that can be achieved with the SKA (http://www.skatelescope.org) and its pathfinders.  The outline for this work is the following: In Sec. 2 we describe how to model non-Gaussian features in LSS bias, and their mass dependence. In Sec. 3 we present the forecasts for future radio surveys, focusing on the improvement of the constraints on $f_{nl}$ when using tracers of different masses with standard angular power spectrum measurements. In Sec. 4, we show how to improve these constraints by eliminating cosmic variance using multi tracer bias. Finally, in Sec. 5 we discuss the results and present our conclusions.  \\section[]{Primordial non-Gaussianity in LSS clustering} \\label{section2} \\subsection{The effect on bias}  In most standard inflationary scenarios, non-Gaussianity in the primordial fluctuation field is characterized by a local feature of the Bardeen gauge invariant potential {$\\Phi$}, which is expressed as: \\begin{equation} \\Phi=\\phi+f_{nl}(\\phi^2-\\langle \\phi \\rangle^2) \\label{eq:1} \\end{equation} where $\\phi$ is a random Gaussian field and $f_{nl}$ is a constant that defines the amplitude of non-Gaussianity. On sub-horizon scales, the Bardeen potential simply becomes the opposite of the gravitational potential. In terms of large scale structure clustering in the Universe, the main consequence of such deviation from a Gaussian field is to increase power on large scales in 2-point statistics such as the tridimensional power spectrum of clustered objects. A scale dependent correction appears in the total halo bias: \\citep{Dalal,b2} \\begin{equation} b_h(M,z)=b_L(M,z)+f_{nl}\\delta_c \\left[(b_L(M,z)-1\\right]\\frac{3 \\Omega_m H_0^2}{c^2 k^2 T(k) D(z)} \\label{eq:2} \\end{equation} The total halo bias thus depends on the non-Gaussianity amplitude $f_{nl}$, the critical overdensity for spherical collapse $\\delta_c$, at redshift z=0, as well as the linear transfer ($T(k)$) and growth ($D(z)$) functions (with $D(0)=1$). This correction is placed upon the usual Gaussian linear bias, which assumes the expression: \\begin{equation} b_L(M,z)=1+\\frac{q \\nu -1}{\\delta_{c}(z)}+\\frac{1}{\\delta_c(z)}\\frac{2p}{1+(q\\nu)^p} \\label{eq:3} \\end{equation} where $\\nu=\\delta_c^2(z)/\\sigma_0^2(M)$ and $\\delta_c(z)= D(z) \\delta_c$. This model was first proposed in \\cite{MoWhite} and \\cite{SethTormen} obtained the values for the parameters which better fit numerical simulations of dark matter collapse and galaxy formation, finding $p=0.3$ and $q=0.75$.   \\subsection{Mass dependence of non-Gaussian effects in the power spectra}  For a given tracer with bias $b_h(z)$, the total 3D halo power spectrum is then simply $P^{3D}_h=b_h^2 P_{\\delta}$ (on large scales), where $P_{\\delta}$ is the underlying dark matter power spectrum. If no redshift information is available to compute the 3D power spectrum, the adequate estimator to use is the angular power spectrum, where the distribution of halos along the line of sight is projected over the 2D field of view. If several tracers are considered, the total statistical information is described by the auto and cross correlation power spectra. The result consists in a set of multipole values $C_l^{i,j}$ and the full computation reads \\citep{Huterer}: \\begin{equation} C_l^{i,j}=\\frac{2}{\\pi}\\int_{k_{min}}^{k_{max}} k^2 P_\\delta(k) W^i_l(k)W^j_l(k)dk \\label{eq:4} \\end{equation} where $W^i_l$ is a window function which accounts for the clustering properties of a given biased tracer $i$ and the angular geometry for a given multipole $l$: \\begin{equation} W^i_l=\\int \\frac{dn}{dz}^iD(z)b_h^i(z)j_l(k r)dz\\;. \\label{eq:5} \\end{equation} Here $dn/dz$ is the angular redshift distribution of sources normalized to unity, $r$ is the comoving radial distance to redshift $z$ and $j_l$ are the spherical Bessel functions of order $l$. \\begin{figure*} \\hspace{-9pt} \\begin{minipage}[b]{8.6cm} \\includegraphics[width=1\\textwidth]{fig1_a.eps} \\end{minipage} \\begin{minipage}[b]{8.6cm} \\includegraphics[width=1\\textwidth]{fig1_b.eps} \\end{minipage} \\captionof{figure}{\\emph{Left:} 3D halo power spectrum at z$=$1 for different values of $f_{nl}$ with an effective halo bias (eq. \\ref{eq:6}) computed in the mass range from $10^{11}h^{-1}M_{\\odot}$ to $10^{12}h^{-1}M_{\\odot}$(lower curves) and from $10^{13}h^{-1}M_{\\odot}$ to $10^{14}h^{-1}M_{\\odot}$ (Upper curves). \\emph{Right:} Ratio between the $P^{3D}_h$ curves presented on the left panel for the same values of $f_{nl}$.} \\label{fig1} \\end{figure*}   As shown in the previous section, the total scale dependent correction introduced by non-Gaussianity in halo clustering ($\\delta_h=b_h\\delta_m$) depends on the redshift and also on the mass of the halo. In the absence of non-Gaussianity, this mass dependence simply translates into an increased amplitude in the power spectrum or correlation function for more massive halos/galaxies when compared to lower mass objects. However, if non-Gaussianity is present, the mass dependence of the linear bias in eq. \\ref{eq:2} will introduce different scale dependence features in clustering 2-point statistics. More massive halos will then be more sensitive to non-Gaussianity than lower mass ones. Fig. 1 shows this effect for different values of $f_{nl}$ and halo mass for the 3D halo power spectrum. One can clearly see the differential effect of using mass bins, with the effect being significant only for $k$ less than 0.02 $h$Mpc$^{-1}$, and having a relative amplitude up to order 1 order of magnitude. Such features offers the possibility to obtain additional information on $f_{nl}$ by constructing observational 2-point statistics from objects corresponding to different halo mass ranges. In order to exploit this, we need to address the issue of differentiating halo masses from an observational point of view. ",
        "conclusions": ""
    },
    "1402/1402.5824_arXiv.txt": {
        "abstract": "Giant flares on soft gamma-ray repeaters that are thought to take place on magnetars release enormous energy in a short time interval. Their power can be explained by catastrophic instabilities occurring in the magnetic field configuration and the subsequent magnetic reconnection. By analogy with the coronal mass ejection (CME) events on the Sun, we develop a theoretical model via an analytic approach for magnetar giant flares. In this model, the rotation and/or displacement of the crust causes the field to twist and deform, leading to flux rope formation in the magnetosphere and energy accumulation in the related configuration. When the energy and helicity stored in the configuration reach a threshold, the system loses its equilibrium, the flux rope is ejected outward in a catastrophic way, and magnetic reconnection helps the catastrophe develop to a plausible eruption. By taking SGR 1806 - 20 as an example, we calculate the free magnetic energy released in such an eruptive process and find that it is more than $10^{47}$ ergs, which is enough to power a giant flare. The released free magnetic energy is converted into radiative energy, kinetic energy and gravitational energy of the flux rope. We calculated the light curves of the eruptive processes for the giant flares of SGR 1806 - 20, SGR 0526-66 and SGR 1900+14, and compared them with the observational data. The calculated light curves are in good agreement with the observed light curves of giant flares. ",
        "introduction": "Soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are believed to be magnetars - a small class of spinning neutron stars with ultra-strong magnetic fields ($\\emph{B}$ $\\geq$ $10^{15}$ G ), which are thought to result from dynamo action during supernova collapse \\citep{DT92,Kel98,TM01,Lyu03,HL06}. The emission from a magnetar is powered by the dissipation of non-potential (current-carrying) magnetic fields in the magnetosphere \\citep{DT92,TD96,TLK02,Lyu06}. Both SGRs and AXPs show quiescent persistent X-ray and repeated soft gamma-ray emissions \\citep{Mel04}. Extremely rarely, an SGR produces a giant flare with enormous energy ($\\backsimeq 10^{44}-10^{47}$ erg) and long burst duration. These exceptionally powerful outbursts begin with a very short ($\\sim$ 0.2~s) spike of $\\gamma$-rays containing most of the flare energy and the spike is followed by a pulsating tail lasting a few hundreds of seconds \\citep{Hur05}.  So far, three SGRs have been reported to produce giant flares \\citep{Maz79,Hur99,Hur05}. They include SGR 0526-66 on 5 March 1979 \\citep{Maz79}, SGR 1900+ 14 on 27 August 1998 \\citep{Hur99,Kou99,Vrb00}, and SGR 1806-20 on 27 December 2004 \\citep{Hur05,Pal05}. The giant flare from SGR 1806- 20 was much more luminous than the other two events \\citep{Hur05,Pal05}. Its initial $\\gamma$-ray spike released an energy of $\\sim 10^{46}$ erg within $\\sim 0.2$~s, and its rising and falling times were $\\tau_{rise}\\leq 1$ ms and $\\tau_{fall}\\approx 65$ ms, respectively. The main spike was followed by a tail with $\\sim 50$ pulsations of high-amplitude at the rotation period (7.56s) of SGR 1806 - 20 \\citep{Hur05,Pal05}.  Although the energy of a magnetar outburst is widely believed to come from the star's magnetic field, details of the physical process in which the magnetic energy is stored and released remain unknown. So far, two models of giant flares of SGRs exist, which depend on the location where the magnetic energy is stored prior to the eruption: one assumes that the energy is stored in the crust of the neutron star (crust model) and the other one assumes that the storage occurs in the magnetosphere (magnetosphere model). In the crust model, a giant flare is caused by a sudden untwisting of the internal magnetic field \\citep{TD95,TD01,Lyu06}. Subsequently, a large and quick rotational displacement on the time-scale of a flare leads to the giant flare. Alternatively, in the magnetosphere model, the magnetic energy is slowly stored in the magnetosphere on time scales much longer than that of the giant flare itself, until the system reaches a critical state at which the equilibrium becomes unstable. Then further evolution in the system occurs and leads to flares, in analogy with solar flares and coronal mass ejections (CMEs) taking place in the solar atmosphere \\citep{Lyu06}. Observations of the giant flare from SGR 1806-20 on 27 December 2004 showed that it lasted a very short rise time, $\\sim$ 0.25~ms \\citep{Pal05}. This time interval of the eruption is short compared to the time-scale required for the crust model \\citep{TD95,Lyu03,Lyu06}. Therefore, at least for the SGR 1806 - 20, the time-scale of the crust model is too long to account for the triggering and early stage evolution of the event.  The magnetosphere model based on an analogy with solar CMEs was proposed by \\citet{Lyu06}. In this model, the magnetic energy released during a giant flare is built up slowly in the magnetosphere, not in the crust of the neutron star. Although the energy storage process is gradual and long, the energy release takes place within a very short timescale in a dynamical fashion \\citep{Lyu06}. Transition from slow to fast evolution constitutes the catastrophe that is similar to what happens in solar flares and CMEs \\citep[see][] {FI91,Iea93}. Moreover, the profile of the light curve of solar-flare/ CME events resembles that of magnetar giant flares. Both of them have an impulsive phase and a tail emission. Similar morphology and characteristics between magnetar giant flares and solar-flare/ CME events indicate the operation of a common physical mechanism. Therefore, solar flares and CMEs give an important prototype context for magnetar giant flares. \\citet{Mas10} constructed a theoretical model for a magnetar giant flares based on the solar-flare/ CMEs model. They described magnetar giant flare using a magnetic reconnection model of a solar flare proposed by \\citet{SY99} while taking account of chromospheric evaporation. In their work, the preflare activity produces a baryon-rich prominence. Then the prominence erupts as a result of magnetic reconnection, and the eruption constitutes the origin of the observed radio-emitting ejecta associated with the giant flare from SGR 1806- 20 \\citep{Tay05,Cam05,Gae05,Mas10}. A giant flare should be induced as the final outcome of prominence eruption accompanied by large-scale field reconfigurations \\citep{Mas10}.  Numerical simulations have also been performed in order to construct an MHD model for a magnetar giant flare. \\citet{Pfr12a,Pfr12b,Pfr13} presented the simulations of evolving strongly twisted magnetic fields in the magnetar magnetosphere. Their results showed that slow shearing of the magnetar crust leads to a series of magnetospheric expansion and reconnection events, corresponding to X-ray flares and bursts. They studied the relationship between the increasing twist and the spindown rate of the star, and concluded that the observed giant flares could be caused by the sudden opening of large amounts of overtwisted magnetic flux, resulting in an abrupt increase in spin period.  Motivated by CME studies, \\citet{Yu11} constructed a  general relativistic model of non-rotating magentars, which simulated how the magnetic field could possess enough energy to overcome the Aly-Sturrock energy constraint and open up. Furthermore, by taking into account the possible flux injections and crust motions, \\citet{Yu12} built a force-free magnetosphere model with a flux rope suspended in the magnetosphere and investigated the catastrophic behavior of the flux rope in a background with multi-polar magnetic field. In this model, a gradual process leads to a sudden release of the magnetosphere energy on a dynamical timescale \\citep{Yu12}. Therefore, the existing catastrophe model for solar eruptions could be a good template for constructing a theoretical model for magnetar giant flares.  The initiation and development of magnetar giant flares have been extensively studied and several models have been suggested \\citep{TD95,TD01,Lyu06, Mas10,Yu11,Yu12}. But the origin and development of these giant flares remains unclear. The detailed physical process of the magnetic energy storage and release is still an open question. In this work, we consider relativistic effects and construct a magnetohydrodynamical (MHD) model for magnetar giant flares in the framework of the CME catastrophe model \\citep{LF00}, then duplicate the dynamical process of the giant flares produced by SGR 1806- 20. We describe our model in next section. Results of calculations and comparisons with observations are in Section 3. Finally, we discuss these results and summarize this work in Section 4. ",
        "conclusions": "We develop a theoretical model via an analytic approach  for a magnetar giant flare in the framework of the solar CME catastrophe model \\citep{LF00}. We considered the physical process that causes the catastrophic loss of equilibrium of a twisted flux rope in the magnetar magnetosphere. The model was constructed according to manifestations of the eruption from the magnetar SGR 1806-20. As happens on the Sun, the free magnetic energy that drives the eruption on the magnetar is slowly stored in the magnetosphere of the magnetar long before the outburst until the system loses mechanical equilibrium, and then is quickly released in the consequent eruptive process, also known as a magnetar giant flare. In our model, the energy driving the eruption comes from the magnetosphere. The motion of footpoints causes the magnetic configuration to lose its equilibrium and release magnetic energy eventually, which is in principle the same as the results of numerical simulations \\citep{Pfr12a,Pfr12b,Pfr13,Yu11,Yu12}.  In the present model, the eruption starts with a loss of equilibrium in the magnetic configuration, which could be triggered either by the change in the background magnetic field, or by the sudden break of a crust piece on the surface of the star where the disrupting magnetic configuration roots in. The evolution in this stage could be purely mechanical, namely no dissipation or magnetic reconnection needs to take place in the magnetosphere. But the dissipation is required in the consequent progress, otherwise the conversion of magnetic energy into radiative and kinetic energy would be stopped, and the evolution in the system ceases without plausible eruptive phenomenon happening \\citep[e.g., see detailed discussions by][]{ FL00,LF00,Lin02,PF02,Lel03,Fel06}.  Our main results deduced from the present analysis are summarized as follow:  1. The system could store free energy of more than $10^{47}$~ergs prior to the eruption, which is enough to drive a giant flare like the one from magnetar SGR 1806-20  observed on 27 December 2004.  2. A combination of the following three factors determines that the disrupting magnetic configuration is highly stretched, and a long magnetically neutral current sheet forms separating two magnetic fields of opposite polarity. These factors are the mechanical property of the loss of equilibrium, the inertia  of the magnetic field that tends to keep the original topological  features in the configuration unchanged without diffusion in the system \\citep[see detailed discussions by][]{FI91,Iea93}, and the fact that the reconnection time-scale is long compared to that of the loss of equilibrium.  3. A long current sheet is usually unstable to perturbations due to various plasma instabilities, such as the tearing mode instability, and the consequent turbulence quickly dissipates the magnetic field such that the stored magnetic energy is rapidly converted into radiative and kinetic energy of the plasma inside the current sheet, as well as the kinetic energy and the gravitational potential energy of the mass in the ejected flux rope. Most of the released magnetic energy turns into radiative energy in the case of the giant flare from SRG 1806-20.  4. Conversion of energy by reconnection mainly occurs at the lower part of the current sheet, and the hottest part of the current sheet is then expected close to the star surface. A large amount of energetic particles and a heat conduction front are also created by reconnection and propagate downward along magnetic field lines. They eventually reach the star surface and may heat the relevant region significantly, which then leads to the formation of a ``fireball\" near the neutron star's surface. Therefore, a ``fireball\" is a straightforward and natural consequence of our model.  5. We calculated the light curve on the basis of our model and compare it with the observational one obtained by the RHESSI $\\gamma$- ray detectors. We note that the calculated light curve consists of a hard spike lasting a half ms and a tail emission last a few $10^{2}$~s, which are consistent with the observations.  6. Our calculations indicate that the magnetic flux preexisting in the flux rope before the eruption was about $7.9\\times 10^{26}$~Mx, and that the magnetic flux brought from the environment around the disrupting magnetic field into the ejecta bubble was around $6.8\\times 10^{26}$~Mx. Therefore, total magnetic flux of more than $1.5\\times 10^{27}$~Mx was sent into interstellar space from the central star by the super eruption from SGR 1806-20.  7. We duplicated our calculations for the giant flares from SGR 0526-66 and SGR 1900+14, respectively, and also found good agreement of our model with observations. Furthermore, the the total magnetic fluxes ejected into interstellar space by the eruption were $6.9\\times 10^{26}$~Mx for SGR 1900+14 and $6.7\\times 10^{26}$~Mx for SGR 0526-66, respectively."
    },
    "1402/1402.1770_arXiv.txt": {
        "abstract": "\\edit{We carry out a systematic study of the X-ray emission from the active nuclei of the $0.02<z<0.7$ 2Jy sample, using \\textit{Chandra} and \\textit{XMM-Newton} observations. We combine our results with those from mid-IR, optical emission line and radio observations, and add them to those of the 3CRR sources. We show that the low-excitation objects in our samples \\redit{show signs} of radiatively inefficient accretion. We study the effect of the jet-related emission on the various luminosities, confirming that it is the main source of soft X-ray emission for our sources. We also find strong correlations between the accretion-related luminosities, and identify several sources whose optical classification is incompatible with their accretion properties. We derive the bolometric and jet kinetic luminosities for the samples and find a difference in the total Eddington rate between the low and high-excitation populations, with the former peaking at $\\sim1$ per cent and the latter at $\\sim20$ per cent Eddington. Our results are consistent with a simple Eddington switch when the effects of environment on radio luminosity and black hole mass calculations are considered. The apparent independence of jet kinetic power and radiative luminosity in the high-excitation population in our plots supports a model in which jet production and radiatively efficient accretion are not strongly correlated in high-excitation objects, though they have a common underlying mechanism.} ",
        "introduction": "Our knowledge of active galactic nuclei (AGN), their observational properties and underlying mechanisms has vastly increased over the last few decades. We now know that these objects are powered through gas accretion onto some of the most supermassive black holes that sit in the centres of most galaxies \\citep[e.g.][]{Magorrian1998}. Radio-loud objects are particularly important to our understanding of AGN, since, despite the fact that they constitute only a small fraction of the overall population, it is during this phase that the impact of the AGN on their surrounding environment (through the production of jets and large-scale outflows and shocks) can be most directly be observed and measured \\citep[e.g.][]{Kraft2003,Cattaneo2009,Croston2011}. Moreover, radio galaxies make up over 30 per cent of the massive galaxy population, and it is likely that all massive galaxies go through a radio-loud phase, as the activity is expected to be cyclical \\citep[e.g.][]{Best2005,Saikia2010}.   It is now commonly accepted that the dominant fuelling mechanism for radio-quiet objects is the accretion of cold gas onto the black hole from a radiatively efficient, geometrically thin, optically thick accretion disk \\citep{Shakura1973}. However, this may not be the case for radio-loud objects. \\citet{Hine1979} noticed the existence of a population of radio-loud objects which lacked the high-excitation optical emission lines traditionally associated with AGN. These so-called low-excitation or weak-line radio galaxies (LERGs or WLRGs) cannot be unified with the rest of the AGN population \\edit{(high excitation galaxies in general, or HEGs, and radio galaxies in particular, or HERGs)}, since their differences are not merely observational or caused by orientation or obscuration. It has been argued that \\edit{LERGs} accrete hot gas \\citep[see e.g.][]{Hardcastle2007b,Janssen2012} in a radiatively inefficient manner, through optically thin, geometrically thick \\edit{accretion flows \\citep[RIAF, see e.g][]{Narayan1995,Quataert2003}}. These objects thus lack the traditional accretion structures (disk and torus) commonly associated with active nuclei \\citep[see e.g.][]{VDWolk2010,FdezOntiveros2012,Mason2012}, and seem to be channeling most of the gravitational energy into the jets, rather than radiative output. This makes them very faint and hard to detect with any non-radio selected surveys.   Current models \\citep[e.g.][]{Bower2006,Croton2006} suggest that the radiatively efficient process may be dominant at high redshifts, \\redit{and to be related to the scaling relation between black hole mass and host galaxy's bulge mass} \\citep[e.g.][]{Silk1998,Heckman2004}. \\edit{Radiatively efficient accretion} may also be the mode involved in the apparent correlation (and delay) between episodes of star-formation and AGN activity in the host galaxies \\citep[e.g][]{Hopkins2012,Ramos2013}. Radiatively inefficient accretion is believed to be more common at low redshifts \\citep{Hardcastle2007b}, and to play a crucial role in the balance between gas cooling and heating, both in the host galaxy and in cluster environments \\citep{McNamara2007,Antognini2012}. These two types of accretion are often called `quasar mode' and `radio mode', which is somewhat misleading, given that there are radiatively efficient AGN with jets and radio lobes. \\redit{This change of a predominant accretion mode with redshift is applicable primarily to the largest galaxies and most massive supermassive black holes (SMBH), since smaller systems evolve differently.}  As pointed out by \\edit{e.g. \\citet{Laing1994,Blundell2001,Rector2001,Chiaberge2002,Hardcastle2009}}, it is important to note that the high/low-excitation division does not directly correlate with the FRI-FRII categories established by \\citet{FR1974}, as is often thought. While most low-excitation objects seem to be FRI, there is a population of bona-fide FRII LERGs, as well as numerous examples of FRI HERGs \\citep[e.g][]{Laing1994}. \\edit{This lack of a clear division} is most likely caused by the complex underlying relation between fuelling, jet generation and environmental interaction. There seems to be a evidence for a difference in the Eddington rate between both populations \\citep[see e.g.][]{Hardcastle2007b,Lin2010,Ho2009,Evans2011,Plotkin2012,Best2012,Russell2012,Mason2012}, with LERGs typically accreting at much lower rates ($<0.1$ Eddington) than HERGs. Estimating the jet kinetic power is also complicated, given that the radio luminosity of a source depends on the environmental density \\citep{Hardcastle2013,Ineson2013} and given the apparent difference in the particle content and/or energy distribution for typical FRI and FRII \\edit{jets and lobes} \\citep[see e.g.][]{Croston2008b,Godfrey2013}.  \\edit{In terms of their optical classification, HERGs are further split into quasars (QSOs), broad-line radio galaxies (BLRGs), and narrow-line radio galaxies (NLRGs), in consistency with the unified models, and in parallel with their radio-quiet counterparts (respectivelly, radio-quiet quasars, type 1 and type 2 radio-quiet AGN). We will use the optical classification for HERGs throughout this work.}  In this paper we analyse the X-ray emission from the 2Jy sample of radio galaxies \\citep{WP1985}, with an approach based on that of \\citet{Hardcastle2006,Hardcastle2009} used on the 3CRR galaxies. X-ray emission is less ambiguous than other wavelengths for an analysis of a sample such as the 2Jy, which contains a variety of populations, in that\\edit{, at these high luminosities, and in the nuclear regions we are considering,} it is unequivocally linked to AGN activity. To fully understand the characteristics of this AGN activity, however, a multiwavelength approach is needed.  From works like those of \\citet{Hardcastle2006,Hardcastle2009}, we do know that LERGs follow the correlation of narrow-line galaxies (NLRGs) between soft X-ray and radio emission \\citep{Hardcastle1999}, reinforcing the hypothesis that in radio-loud objects this X-ray component originates in the jet. One of the crucial points we aim to investigate in this paper is the dissimilarity between the NLRG and LERG populations.  Our aim is to study the correlations between the luminosities of the sources at different wavelengths, to \\edit{link the emission produced in regions at various distances from the central black hole: from the disk and corona to the torus, the jet and the lobes. In doing so, we will investigate how accretion translates into radiative and kinetic output across the whole radio-loud population.}  \\edit{While many of the sources in the 3CRR catalogue have been observed in great detail, the multiwavelength coverage is not uniform, and the sample is not statistically complete in the X-rays, being more complete for redshifts $<$0.5. The observations of the 2Jy sample, however, were taken with the explicit purpose of providing \\redit{comparable measurements} for all the objects in the sample. This consistency provides us with the opportunity to test whether the conclusions reached by Hardcastle et al. can be extrapolated to all radio-loud AGN or are related to the biased redshift distribution of the 3CRR sources.}  \\edit{Although it is well known that some of the physical mechanisms involved in radio-loud emission in AGN are similar to those found in X-ray binaries \\citep[see e.g. the review by][]{Kording2006}, some caution must be applied, since there are also dissimilarities in the timescales and fuelling processes involved. In this work we will focus only on AGN, and the possible impact of our results may have on understanding their observational properties, calssification, accretion mode and the influence on their hosts.}  For this paper we have used a concordance cosmology with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": ""
    },
    "1402/1402.4363_arXiv.txt": {
        "abstract": "Gravitationally coupled scalar fields $\\phi $, distinguished by the choice of an effective self-interaction potential $V(\\phi )$, simulating a temporarily non-vanishing cosmological term, can generate both inflation and late time acceleration. In scalar field cosmological models the evolution of the Hubble function is determined, in terms of the interaction potential, by a Riccati type equation. In the present work we investigate scalar field cosmological models that can be obtained as solutions of the Riccati evolution equation for the Hubble function. Four exact integrability cases of the field equations are presented, representing classes of general solutions of the Riccati evolution equation. The solutions correspond to cosmological models in which the Hubble function is proportional to the scalar field potential plus a linearly decreasing function of time, models with the time variation of the  scalar field potential proportional to the potential minus its square, models in which the potential is the sum of an arbitrary function and the square of the function integral, and models in which the potential is the sum of an arbitrary function and the derivative of its square root, respectively. The cosmological properties of all models are investigated in detail, and it is shown that they can describe the inflationary or the late accelerating phase in the evolution of the Universe. \\\\ \\\\ {\\bf keywords}: Riccati equation; isotropic scalar field; cosmology ",
        "introduction": "The standard model of cosmology is remarkably successful in accounting for the observed features of the Universe. However, there remain a number of fundamental open questions at the foundation of the standard model. In particular, we lack a fundamental understanding of the acceleration of the late universe \\cite{latetime}. In fact, the standard model of cosmology has favoured dark energy models, involving time-dependent scalar fields, as fundamental candidates responsible for the cosmic expansion. Indeed, scalar fields naturally arise in particle physics, including string theory, and in addition to this, the underlying dynamics in inflationary models depend essentially on a single scalar field, with the inflaton rolling in some underlying potential \\cite{1}. A plethora of candidates exist for dark energy, and we refer the reader to \\cite{Copeland:2006wr} for a review.  In a cosmological context, the mathematical properties of the Friedmann-Robertson-Walker (FRW) models with a scalar field as a matter source have also been extensively investigated. For instance, in \\cite% {Mus,Sal} a simple way of reducing the system of the gravitational field equations to a first order equation was proposed, namely, to the Hamilton-Jacobi-like equation for the Hubble parameter $H$ considered as a function of the scalar field $\\phi $. In fact, a number of integrable one--scalar spatially flat cosmologies, which play a natural role in the inflationary scenarios, were studied in \\cite{Fre1, Fre2}. Recently, a general method for the study of scalar field cosmologies, based on the reduction of the Klein-Gordon equation to a first order non-linear differential equation, was proposed in \\cite{HLM}.  It is interesting to note that the possibility of describing the cosmological dynamics for a barotropic fluid in terms of a Riccati equation was discussed in \\cite{Far}. For a cosmological fluid satisfying an equation of state of the form $p=(\\gamma -1)\\rho $, the Friedmann equations give for the scale factor the evolution equation in the conformal time $\\eta $ is written as $a^{\\prime \\prime }/a+(c-1)\\left( a^{\\prime }/a\\right) ^{2}+ck=0$% , where $c=3\\gamma /2-1$, and $k=\\pm 1$, and a prime denotes the derivative with respect to $\\eta $. By introducing the transformation $u=a^{\\prime }/a$% , we obtain for $u$ the Riccati type equation $u^{\\prime }+cu^{2}+ck=0$, which is easily integrable. The Riccati equation based study of the different properties of the isotropic FRW type cosmological models was performed in \\cite{Rosu,10b,10c, 10d,10e}. In fact, the integrability conditions of the Riccati equation obtained in \\cite{MHR,MHR2} allow the integration of the structure equations of isotropic general relativistic compact objects \\cite% {MHS} in the context of general relativity. Furthermore, the applications of the Riccati equation to stellar and cosmological models have been extensively discussed in the literature, and we refer the reader to \\cite% {MH6,MH3,MH4,MH5,U1}. Very recently, using the Chiellini type integrability condition for the generalized first kind Abel differential equation, consequently, the new class of exact analytical solution of the Riccati type equation has been obtained in \\cite{U2}.  Therefore, the theoretical investigation of scalar field models is an essential task in cosmology. It is the purpose of the present paper to explore alternative approach in solving the cosmological gravitational field equations in the presence of self-interacting scalar fields, based on the mathematical analysis of the Riccati type equation that gives the Hubble function in terms of the scalar field potential. We present several cases of integrability of this equation, corresponding to specific forms of the scalar field potential, or of the Hubble function. The physical properties of the obtained solutions are analyzed in detail.  The present paper is organized as follows. The basic Riccati evolution equation for scalar field cosmologies with an arbitrary self-interaction potential is derived in Section \\ref{sect2}, and four classes of exact scalar field solutions are obtained in Section \\ref{sect3}. We discuss and conclude our results in Section \\ref{sect7}. Throughout this paper, we use natural units $c=8\\pi G=\\hbar =1$, and adopt as our signature for the metric $\\left( +1,-1,-1,-1\\right) $. ",
        "conclusions": "\\label{sect7}  In the present paper, we have shown that the time evolution and dynamics of the Hubble function in scalar field cosmologies can be formulated in terms of a simple first order Riccati type equation, with the cosmological dynamics entirely determined by the time variation of the scalar field potential. This equation immediately leads to the identification of some classes of scalar field potentials for which the field equations can be solved exactly, and it allows the formulation of very general integrability conditions. We have obtained the complete solution of the gravitational field equations describing the time evolution of the flat FRW Universe in the presence of the scalar field $\\phi \\left( t\\right) $ for four functional forms of the self-interaction potential $V$. The first two solutions are obtained for fixed forms of the scalar field potential, while in the last two solutions the form of the potential is arbitrary, and determined by a general integrability condition. The integrability conditions determine the allowed form of the scalar field self-interaction potential in terms of some arbitrary time dependent functions $f_{1}\\left( t\\right) $ and $f_{2}\\left( t\\right) $, respectively, thus leading to the possibility of constructing very general solutions of the field equations, and to reconstruct easily the Hubble function, once the evolution of the potential is given.  In conclusion, we have obtained several exact solutions of the gravitational field equations in the presence of a scalar field. In order to obtain a deeper physical understanding of the solutions  comparisons with the observational data are necessary. Work under these lines is presently underway, and the results will be presented in a future publication."
    },
    "1402/1402.3296_arXiv.txt": {
        "abstract": "We make use of a semi-analytical model of galaxy formation to investigate the origin of the observed correlation between [$\\alpha /{\\rm Fe}$] abundance ratios and stellar mass in elliptical galaxies. We implement a new galaxy-wide stellar initial mass function (Top Heavy Integrated Galaxy Initial Mass Function, TH-IGIMF) in the semi-analytic model \\sag~and evaluate its impact on the chemical evolution of galaxies. The SFR-dependence of the slope of the TH-IGIMF is found to be key to reproducing the correct [$\\alpha /{\\rm Fe}$]-stellar mass relation. Massive galaxies reach higher [$\\alpha /{\\rm Fe}$] abundance ratios because they are characterized by more top-heavy IMFs as a result of their higher SFR. As a consequence of our analysis, the value of the minimum embedded star cluster mass and of the slope of the embedded cluster mass function, which are free parameters involved in the TH-IGIMF theory, are found to be as low as $5\\,\\Msun$ and $2$, respectively. A mild downsizing trend is present for galaxies generated assuming either a universal IMF or a variable TH-IGIMF. We find that, regardless of galaxy mass, older galaxies (with formation redshifts $\\gtrsim 2$) are formed in shorter time-scales ($\\lesssim 2\\,{\\rm Gyr}$), thus achieving larger [$\\alpha /{\\rm Fe}$] values. Hence, the time-scale of galaxy formation alone cannot explain the slope of the [$\\alpha /{\\rm Fe}$]-galaxy mass relation, but is responsible for the big dispersion  of [$\\alpha /{\\rm Fe}$] abundance ratios at fixed stellar mass. We further test the hyphothesis of a TH-IGIMF in elliptical galaxies by looking into mass-to-light ratios, and luminosity functions. Models with a TH-IGIMF are also favoured by these constraints. In particular, mass-to-light ratios agree with observed values for massive galaxies while being overpredicted for less massive ones; this overprediction is present regardless of the IMF considered. \\\\ ",
        "introduction": "\\label{sec:intro}  The chemical signatures imprinted in stellar populations are key features in the understanding of the formation and evolution of galaxies. The alpha-element abundance relative to iron has been recognized for a long time as an indicator of the formation time-scale of stellar populations.  It is well known that the alpha-elements are produced by alpha capture exclusively in core-collapse supernovae (SNe CC), whose progenitors have lifetimes of tens of Megayears, while Fe and Fe-peak elements are mainly the yields of type Ia supernovae (SNe Ia), whose progenitors are long lived \\citep[$0.5 - 3 \\,{\\rm Gyr}$,][]{Greggio05}. A composite stellar population with an excess of alpha-elements relative to iron is therefore assumed to form in a short period of time ($\\lesssim 1 \\,{\\rm Gyr}$) \\citep{Matteucci94}, given that the bulk of stars in the system form before the interstellar medium (ISM) is contamined with the SNe Ia yields. The early works of \\citet{Tinsley76, Tinsley79} already considered this abundance ratio in order to explain the formation time-scales of the solar neighbourhood and the Milky Way stellar halo. This argument has been also considered to explain the trend observed among elliptical galaxies in which more massive ellipticals (with higher velocity dispersions) have systematically higher alpha-elements to iron ratios \\citep[e.g.][]{Worthey1992}. \\citet{Faber1973} first observed that some absorption features related to the abundance of Mg increase monotonically with increasing luminosity. Since then, a large amount of work has been done measuring alpha-element abundances to study this trend in all sort of environments \\citep[e.g.][]{Worthey1992, Trager2000, Kuntschner2000, Bernardi2006, SanchezBlazquez2006, Spolaor2010, Thomas2010}.  In the context of the $\\Lambda$-CDM paradigm, the observed trend of $[\\alpha/{\\rm Fe}]$ as a function of mass could be considered troublesome. In this framework, the growth of structure is hierarchical. Dark matter (DM) haloes form via dissipationless gravitational collapse \\citep{White-Rees1978} and grow in a bottom-up fashion, with more massive haloes taking longer periods of time to build-up their masses. Considering that galaxies form in the centre of DM haloes, a priori it is not expected that more massive galaxies have higher abundances of $\\alpha$-elements relative to iron.  Nevertheless, the existence of baryonic processes like gas cooling, star formation (SF) and feedback from SNe and active galactic nuclei (AGN) can detach the actual formation of the galaxy from the assembly of its parent DM halo. This ``anti-hierarchical\" behavior, usually called {\\it downsizing}, manifests in different ways in galaxies: the characteristic mass of star forming galaxies increases with redshift \\citep[e.g.][]{Cowie-Barger2008}, the evolution of the stellar mass functions seems to indicate that less massive galaxies assembly their mass at later times \\citep[e.g.][]{PerezGonzalez2008}, more massive galaxies have older stellar populations than less massive ones \\citep[e.g.][]{Gallazzi2005}. \\citet{Fontanot2009} classify the different observational manifestations of downsizing, suggesting the possibility that not all these phenomena are related to the same underlying physical process. In particular, they coined the term {\\it{chemo-archaeological downsizing}} to refer to the trend of abundance ratios with stellar mass shown by elliptical galaxies.  The pointed tension between observations and the hierarchical paradigm has been addressed by different groups using analytic and semi-analytic models of galaxy formation (SAMs). \\citet{Thomas1999} was the first in modelling the chemical enrichment of elliptical galaxies within the framework of hierarchical galaxy formation, but considering a closed box model for each galaxy and ignoring the merger history of galaxies by using pre-calculated star formation histories. He fails to reproduce the observed trend of $[\\alpha/{\\rm Fe}]$ versus velocity dispersion. \\citet{Nagashima+2005b} study the metal enrichment of elliptical galaxies using a semi-analytic model with merger trees. They find that a top heavy initial mass function (IMF) in starbursts is needed in order to reproduce the observed  $[\\alpha/{\\rm Fe}]$ abundance ratio in $L_{\\star}$ ellipticals, but none of their models reproduce the observed trend of this ratio with velocity dispersion, $\\sigma$.  \\citet{Pipino+2009} focussed on the $[\\alpha/{\\rm Fe}]$-stellar mass relation in elliptical galaxies by including the enrichment of supernovae type Ia and II, and of low- and intermediate- mass stars, as well as the corresponding delay times for SNe Ia in the semi-analytic model of galaxy formation GalICs. They find a better agreement with the observed $[\\alpha/{\\rm Fe}]$-$\\sigma$ relation when taking into account AGN feedback in their models, although still marginal. \\citet{Calura2009} implement a chemical model on top of the star formation histories extracted from a semi-analytic model, and evaluate the impact on the $[\\alpha/{\\rm Fe}]$-$\\sigma$ relation of the AGN feedback and an ad-hoc variable IMF that becomes top heavy in starbursts, finding that they are both needed to obtain a slope in the relation similar to the observed one. This was the first attempt to include a variable IMF in a semi-analytic model to study the $[\\alpha/{\\rm Fe}]$ ratios in ellipticals. A similar assumption about the IMF has been considered by Baugh et al. (2005) to explain the number counts of submillimiter galaxies at high redshift. Later on, \\citet{Calura2011} mention the importance of the harassment of satellite galaxies at high redshift as a trigger of SF in the progenitors of massive galaxies that form stars enhanced in $\\alpha$-elements, contributing to build up the relation. \\citet{Arrigoni2010} also analyse the $[\\alpha/{\\rm Fe}]$-mass relation by using the semi-analytic model of \\citet{Somerville2008} that includes AGN feedback, in which they implement a chemical enrichment model that involves SNe Ia and SNe II yields and the corresponding delay times for the SNe Ia explosions. They find a good agreement with observations by combining a shallower slope for the IMF and a lower SNe Ia fraction. All the results obtained from semi-analytic models of galaxy formation and evolution suggest that the time-scale argument, which compares the time-scales of SF and metal ejection from SNe, might be insufficient to explain the build-up of the $[\\alpha/{\\rm Fe}]$-mass relation of early type galaxies, supporting the importance of the role of the IMF.  Numerous works are being published that show the need to explore alternative and variable stellar initial mass functions. The stellar IMF has direct consequences on the evolution of galaxies. It determines the ratio of massive stars to low and intermediate-mass stars, imprinting the chemical abundance ratios of stellar populations in galaxies \\citep{Lucatello2005, Koppen2007}, and establishing the amount of reheated mass deposited in the ISM by SNe explosions. It also determines the mass-to-light ratio (M/L) of galaxies, given that stars with different initial stellar masses contribute with very different amounts of light along their lives. In the last decade, a formalism for a galaxy-wide stellar IMF has been developed \\citep[see][ or Subsection \\ref{sec:varIMF} of this paper]{Kroupa2012}, where the instantaneous star formation rate (SFR) of galaxies determines the final shape of the integrated IMF. \\citet{Recchi+2009} evaluate the influence on the [$\\alpha$/Fe]-$\\sigma$ relation of the galaxy-wide integrated galactic initial mass function \\citep[IGIMF,][]{Kroupa2003} implemented in an analytical model of chemical enrichment of galaxies. They find that the downsizing effect (shorter duration of the SF in larger galaxies) must be present to reproduce the observed trend, although it must be milder than the one infered by \\citet{Thomas2005} on the basis of the observational [$\\alpha$/Fe]-$\\sigma$ relation. Several simplifying assumptions have been adopted in their work. Moreover, they do not include the effect of a top heavy version of the IGIMF (TH-IGIMF) \\citep{Weidner2011, Weidner2013a}. This particular IMF has been tested by \\citet{Fontanot2014} in the semi-analytic model MORGANA evaluating its impact on the stellar mass functions and star formation rates. These results are compared with those obtained from other IMF with different dependencies on galaxy properties and environment, finding that the predicted galaxy properties are mainly affected by an increasing Top-Heavy IMF at increasing star formation rates.   The aim of this work is to address the chemical evolution of galaxies in the framework of the $\\Lambda$-CDM paradigm, tackling the problem of $\\alpha$-element enrichment in elliptical galaxies. We evaluate the impact of the new theory of TH-IGIMF, never considered so far for this particular issue. We also analyse the influence of the TH-IGIMF in the luminosity function and the mass-to-light ratios of elliptical galaxies analyzing the tilt of the fundamental plane (M/L vs M). For this purpose, we use the state of the art semi-analityc model \\sag~(acronym for Semi-Analytic Galaxies) that couples the most relevant ingredients of chemical evolution together with the physics involved in the formation and evolution of galaxies (\\citealt{Cora06}; \\citealt*{lcp08}; \\citealt{Tecce2010}).  This work is organized as follows. In Section \\ref{sec:HybridModel}, we present the model of galaxy formation used that combines a cosmological dark-matter only simulation and the semi-analytic model of galaxy formation \\sag, giving a brief description of the latter with emphasis on the formation of elliptical galaxies, the chemical aspects and the new implementation of extended bursts of star formation (Subsection \\ref{sec:SAG}); Subsection \\ref{sec:varIMF} is devoted to the detailed description of the implementation of the TH-IGIMF. In Section \\ref{sec:CalibrationSAG}, we present a summary of the calibration process involved in this work and show the first results of the [$\\alpha/{\\rm Fe}$]-stellar mass relation with a universal Salpeter IMF (\\ref{sec:AlphaSalp}). In Section \\ref{sec:AlphaFeTHIGIMF}, we analyse the impact of a variable IMF on the developement of the [$\\alpha/{\\rm Fe}$]-stellar mass relation, discussing the influence of the various time-scales involved in the modelling of formation of elliptical galaxies, such as starburst timescales and SNe delay times  (Subsection \\ref{sec:starbTimeScales}). In Subsection \\ref{sec:masstolight} we test the implementation of the variable IMF by means of the resultant luminosity functions and mass-to-light ratios. In Section \\ref{sec:timescales} we study in detail the global formation timescales of galaxies and analize their impact on the development of the [$\\alpha/{\\rm Fe}$]-stellar mass relation and the existence of a downsizing pattern. In Section \\ref{discussion}, we mention previous investigations in the field putting in context our own results, with the perspective of future research work related to the main topic of this paper. We present our conclusion in Section \\ref{conclusions}. Finally, we add complementary information in appendix \\ref{App:AppendixA}, where we detail the physical recipies implemented in the model and the link with the free parameters involved in the calibration procedure that allows to obtain the best-fitting values.  \\section[]{Model of galaxy formation} \\label{sec:HybridModel} In the present work, we use a model of galaxy formation that combines a dissipationless cosmological {\\em N}-body simulation with a semi-analytic model of galaxy formation and evolution. In the following, we describe the features of the simulation used and those aspects of the semi-analytic code that are relevant for the present study.   \\subsection{{\\em N}-body simulation}  The DM only {\\em N}-body simulation used for this study considers the standard $\\Lambda$-CDM scenario, characterized by cosmological parameters $\\Omega_{\\rm m}=0.28$, $\\Omega_{\\rm b}=0.046$, $\\Omega_{\\Lambda}=0.72$, $h=0.7$, $n=0.96$, $\\sigma_{8}=0.82$, according to the WMAP7 cosmology \\citep{jarosik2011}. The simulation was run with \\gadget~\\citep{springel_gadget2_2005} using $640^3$ particles in a cubic box of comoving sidelength $L=150\\,h^{-1}{\\rm Mpc}$. The initial conditions were generated using {\\small GRAFIC2} \\citep{bertschinger_grafic2_2001}. The simulation was evolved from $z_{\\rm ini}=61.2$ to the present epoch, storing 100 outputs equally spaced in $\\log_{10}(a)$ between $z=20$ and $z=0$ \\citep{benson_trees_2012}.  DM haloes were identified using a friends-of-friends (\\fof) algorithm. Then, the application of \\subfind~algorithm \\citep{springel_subfind_2001} allows to select self-bound substructures (subhaloes) having at least 10 particles, with a mass of $1\\times 10^{9}\\,{h^{-1}}\\,\\Msun$ each. DM haloes and subhaloes have masses in the range $1\\times 10^{10} - 7.8 \\times 10^{14}\\,{h^{-1}} \\,\\Msun$.  \\subsection[]{Semi-analytic model of galaxy formation} \\label{sec:SAG}  We use the semi-analytic model {\\small SAG} described in the works of \\citet{Cora06}, \\citet{lcp08} and \\citet{Tecce2010}, which is based on the version of the Munich semi-analytic model presented by \\citet{springel_subfind_2001}. In this section, we describe briefly the main features of the model, giving more details of aspects closely related with the aim of this research work, that is, the chemical implementation and the channels of bulge formation that give rise to the population of elliptical galaxies. The time-scales involved in the formation of this stellar component are crucial in the understanding of the development of the [$\\alpha$/Fe]-stellar mass relation. Several modifications have to be included for a proper treatment of this aspect, which are also described in detail.  \\subsubsection{Main features of \\sag} \\label{sec:MainFeatures}  The semi-analytic model {\\small SAG} uses the subhalo merger trees provided by the underlying simulation to generate the galaxy population. Cold gas is settled in the centres of DM subhaloes as a result of the radiative cooling suffered by the hot gas contained within them. It is assumed that the cold gas disc has an exponential density profile \\citep{Tecce2010}. When the density of cold gas becomes high enough, SF in the disc is triggered according to the conditions given by \\citet{Croton2006} which are considered in the current version of \\sag. We refer to this mode of SF as quiescent SF. One aspect that changed with respect to previous versions of \\sag~is the estimation of the gas cooling rate; we now consider the total radiated power per chemical element given by \\citet{foster2012}\\footnote{ These modifications were already taken into account in the version of \\sag~considered to evaluate the capability of the calibration method used to tune the free parameters of the model \\citep{Ruiz2014}.}. The only galaxies for which infall of cooling gas from the intergalactic medium takes place are central galaxies. These galaxies, also named {\\it type 0}, reside within the most massive subhalo within a \\fof~halo. Satellite galaxies are classified in two types, those residing in smaller subhaloes of the same \\fof~halo ({\\it type 1}), and those which DM subhalo are below the resolution limit as a result of mass loss caused by tidal forces ({\\it type 2}). Galaxies of the latter type merge with the central galaxy of their host subhalo after a dynamical friction time-scale \\citep{BinneyTremaine87}. When a galaxy becomes a satellite of either type, it suffers strangulation, that is, all of its hot gas halo is removed and transferred to the hot gas component of the corresponding type~0 galaxy. From a given star formation event and an IMF adopted, we can estimate the number of SNe CC that contribute to energy feedback. Effects of feedback from active galactic nuclei (AGN) are also included \\citep{lcp08}. The energy injection from both sources regulates the star formation through the transfer of gas and metals from the cold to the hot gas phase.   \\subsubsection{Formation of elliptical galaxies} \\label{sec:ellip_form}  The morphological classification of galaxies generated by the model is based on the ratio between bulge and total {\\em r}-band luminosity, B/T. Elliptical galaxies are those characterized by a ratio ${\\rm B/T}>{{(\\rm B/T)}_{\\rm thresh}}$, where the parameter ${(\\rm B/T)}_{\\rm thresh}$ may adopt values in the range $\\sim 0.7-0.85$. Thus, the formation of elliptical galaxies is directly connected with bulge formation, which in \\sag~takes place through global disk instabilities and two different kinds of galactic mergers.  When a galaxy merger occurs, the stellar mass ratio between the satellite galaxy and the central galaxy, $M_{\\rm sat}/M_{\\rm cen}$, is evaluated. If $M_{\\rm sat}/M_{\\rm cen}>0.3$, then the merger is considered as a major one. In this case, all the gas in the remnant galaxy is consumed in a starburst contributing to the bulge formation, and the stellar disc is completely relaxed and transferred to the bulge. The triggering of a starburst in a minor merger ($M_{\\rm sat}/M_{\\rm cen}\\leq 0.3$) will depend on the fraction of cold gas present in the disc of the central galaxy, as implemented by \\citet{lcp08} following the work of \\citet{malbon2007}. If the ratio between the cold gas and the disc mass of the central galaxy, $M_\\text{cold,cen} / M_\\text{disc,cen}$, is larger than a fixed parameter $f_\\text{burst}=0.6$, then the perturbation introduced by the merging satellite drives all the cold gas from both galaxies into the bulge component, where it is consumed in a starburst; the disc mass is given by the sum of the cold gas and the stars formed through quiescent SF. However, if the satellite is much less massive than the central galaxy ($M_{\\rm sat}/M_{\\rm cen}\\leq$~0.05), no burst occurs. In minor mergers, only the stars of the merging satellite are transferred to the bulge component of the central galaxy.  The other channel that contributes to the bulge formation is the global disk instability. Some configurations of galactic discs do not remain stable with time. When a galactic disc is sufficiently massive that its self-gravity is dominant, it becomes unstable. This condition is expressed in the model through the Efstathiou-Lake-Negroponte \\citep{Efstathiou+1982} criterium, that is, the stability to bar formation is lost when \\begin{equation} \\epsilon_{\\rm d} \\equiv \\frac{V_{\\rm disc} }{ (G M_{\\rm disc} / R_{\\rm disc})^{1/2}} \\le \\epsilon_{\\rm thresh}, \\label{eq:diskinstab} \\end{equation} where $M_{\\rm disc}$ is the mass of the disc (cold gas plus stars), $R_{\\rm disc}$ is the disc scale radius, and $V_{\\rm disc}$ is the circular velocity of the disc. For the latter, we use the velocity where the rotation curve flattens, which we approximate by the velocity calculated at $\\sim 3\\,R_{\\rm disc}$ (see \\citet{Tecce2010} for details concerning disc features). The free parameter $\\epsilon_{\\rm thresh}$ has values close to unity; in our model, we adopt $\\epsilon_{\\rm thresh}=1$. We consider an additional free parameter that takes into account the influence of a perturbing galaxy that effectively triggers the disc instability; this is modelled by computing the mean separation between galaxies in a \\fof~group. When the mean separation is smaller than $f_\\text{pert}\\,R_{\\rm disc}$, being $f_\\text{pert}$ a free parameter, we consider that a neighbouring galaxy perturbs the unstable disc. As a consequence of this, all the stars and cold gas in that disc are transferred to the bulge component, and all the gas present is consumed in a starburst.   \\subsubsection{Chemical enrichment model} \\label{sec:chem_model}   In our chemical model, we follow the production of ten chemical elements (H, $^4$He, $^{12}$C, $^{14}$N, $^{16}$O, $^{20}$Ne, $^{24}$Mg, $^{28}$Si, $^{32}$S, $^{40}$Ca) generated by stars in different mass ranges, from low- and intermediate-mass stars to quasi massive and massive stars. This version of \\sag~is characterized by a new set of stellar yields. We use the best combination of stellar yields reported by \\citet{Romano2010}, selected to be in accordance with the large number of constraints for the Milky Way. For low and intermediate-mass stars (LIMS), in the mass interval $1-8 \\Msun$, we use the yields of \\citet{Karakas2010}. For the mass loss of pre-supernova stars (He and CNO elements), we use the yields computed by the Geneva group \\citep{Hirschi2005}, and for the yields of the explosive nucleosynthesis due to SNe CC, we use the results of \\citet{Kobayashi2006}. In all cases, we have adopted the total ejected mass of solar metallicity models, for which the solar abundances of \\citet{AndersGrevesse89} are assumed, with a solar composition $Z_{\\odot} = 0.02$.  Ejecta from SNe Ia are also included, which are characterized by high iron production ($\\sim 0.6\\, {\\rm M}_{\\odot}$); we consider the nucleosynthesis prescriptions from the updated model W7 by \\citet{Iwamoto99}. The SNe Ia rates are estimated using the single degenerate model, in which a SN Ia occurs by carbon deflagration in C-O white dwarfs in binary systems whose components have masses between $0.8$ and $8\\,\\Msun$ \\citep{Greggio83}; we implement the formalism presented by \\citet{Lia2002} (see their Section 4.2.2). The fraction of these binary systems, $A_{\\rm bin}$, is one of the free parameters of \\sag. We take into account the return time-scale of mass losses and ejecta from all sources considered. For this purpose, we use the stellar lifetimes given by \\citet{Padovani93} who use results from \\citet{Matteucci86} for masses over $6.6\\,\\Msun$, and from \\citet{Renzini86} for masses below this limit. This aspect becomes especially relevant for this work because of the different delay times that characterize different types of SNe, which affect the abundance of $\\alpha$-elements relative to iron. For a detailed description of the implementation in \\sag~of the chemical enrichment of the different baryonic components, we refer the reader to \\citet{Cora06}.  \\subsubsection{Extended starbursts} \\label{sec:ExtendedBurstsInSAG}  In previous versions of the model, the formation of stars in the bulge occurs through starbursts that consume the available cold gas in a single step \\footnote{Differential equations in \\sag~are integrated in time-steps of equal size used to subdivide the intervals between simulation outputs.}. We now consider that starbursts are characterized by a given time-scale during which the cold gas driven to the galactic center is consumed gradually in several steps. Thus, this cold gas has the possibility of being progressively contaminated by the stars formed in the bulge. This way, the $[\\alpha/{\\rm Fe}]$ abundance ratio of successive generations of stars in the bulge is not only determined by the abundances of the disc cold gas at the moment in which the starburst is triggered, but is also modified by the relative contribution of different types of SNe, which depends on the relation between the time-scale of the duration of the starburst and the lifetime of SNe progenitors.  The cold gas that will be eventually converted in bulge stars is referred to as bulge cold gas, in order to differenciate it from the disc cold gas. The amount of bulge cold gas available allows us to make a first estimation of the mass of stars that will be formed in the bulge. Then, we compute the time-scale in which they would form, which is chosen to be the dynamical time considered in quiescent SF as a first approximation, that is, the one defined as the ratio between the scale radius of the exponential profile that characterizes the galactic disc and its circular velocity. This choice seems reasonable since a burst normally takes place via a bar embedded in the galaxy disc. Thus, from the ratio of the bulge cold gas mass and the time-scale estimated for the duration of the starburst, we know the mass of stars that will be formed in each time step. For simplicity, SF throughout the extended starburst process occurs at a constant rate. Other works in the literature that have taken into account the starbursts time-scales consider that the SFR in a burst decays exponentially with time after the burst is triggered \\citep{Granato2000, Lacey2008}. We justify our assumption considering that the bar formed during the instability, triggered either by a galaxy merger or a disk instability, favours the gradual feeding of the bulge cold gas reservoir by the disc cold gas that is being driven to the galactic center \\citep[e.g.][]{KormendyKennicutt2004}, keeping constant the conditions that give rise to SF. Note that the presence of a bar during the process of bulge formation is only implicit in the assumption of constant SFR during starbursts. Modelling bar formation is not an easy task \\citep[see][]{KormendyKennicutt2004, Athanassoula2013}, and is beyond the scope of this paper.  \\subsection{Variable IMF} \\label{sec:varIMF}  The stellar IMF has a great influence on the chemical enrichment of galaxies since it defines the amount of stars formed in each stellar mass interval for each star formation event and, consequently, impacts on the amount of metals returned to the ISM, where new stars will eventually form. The IMF also determines the number of SNe CC involved in the estimation of the reheated mass that is transferred from the cold to the hot phase during SN feedback. Today, there is some consensus that in a simple stellar population the slope of the IMF above $1\\,\\Msun$ is not strictly different from that of Salpeter IMF ($\\alpha \\equiv 1+x = 2.35$); stars with masses over $1\\,\\Msun$ are the major contributors to chemical enrichment \\citep{Portinari98b}. Several attempts have been made in measuring this important distribution function in the solar neighbourhood \\citep{Salpeter55, Chabrier2003} or directly in star clusters \\citep{Kroupa2001, Kroupa2002}. The latter is represented by a multi-component power-law IMF given by \\begin{equation} \\xi(m) = k \\left\\{\\begin{array}{ll} k'\\,\\left(\\,\\frac{m}{m_{\\rm H}} \\right)^{-\\alpha_{0}}&\\hspace{-0.25cm},{m_{\\rm low}} \\le m/{M}_\\odot < {m_{\\rm H}}\\\\ \\left(\\frac{m}{m_{\\rm H}} \\right)^{-\\alpha_{1}}&\\hspace{-0.25cm},m_{\\rm H} \\le m/{M}_\\odot < m_0,\\\\ \\left(\\frac{m_{0}}{m_{\\rm H}} \\right)^{-\\alpha_{1}} \\left(\\frac{m}{m_{0}} \\right)^{-\\alpha_{2}}&\\hspace{-0.25cm},m_0 \\le m/{M}_\\odot < m_1,\\\\ \\left(\\frac{m_{0}}{m_{\\rm H}} \\right)^{-\\alpha_{1}} \\left(\\frac{m_{1}}{m_{0}} \\right)^{-\\alpha_{2}} \\left(\\frac{m}{m_{1}} \\right)^{-\\alpha_{3}}&\\hspace{-0.25cm},m_1 \\le m/{M}_\\odot \\le m_\\mathrm{max},\\\\ \\end{array} \\right. \\label{eq:4pow} \\end{equation} \\noindent with $m_{\\rm low}=0.01$, $m_{\\rm H}=0.08$, $m_{0}=0.5$ and $m_{1}=1.0$, and exponents $\\alpha_0=0.3$, $\\alpha_1=1.3$, $\\alpha_2=\\alpha_3=2.35$ . The normalization constant $k$ contains the desired scaling of the IMF (for example, to the total mass of the system). Note that  brown dwarfs are a separate population and contribute about 1.5 per cent by mass only. Therefore, they have a different normalisation factor, $k'$ \\citep[$k' \\sim 1/3$,][]{ThiesKroupa2007, ThiesKroupa2008}. In the case of stellar systems with composite stellar populations, other aspects must be considered.  A galaxy-wide IMF may be represented by an integrated galactic initial mass function (IGIMF) proposed by \\citet{Kroupa2003}, whose theory was developed in the last decade \\citep[see][ for a review]{Kroupa2012}. This formalism states that star formation takes place exclusively in star clusters \\citep{LadaLada2003}. This way, the stellar populations of each galaxy are composed by stars in surviving and dissolved star clusters. Within each star cluster, stars form with the stellar IMF given by eq. \\ref{eq:4pow}. The embedded star clusters in the galaxy also follow an initial distribution of masses of the form $\\xi_{\\rm ecl}(M_{\\rm ecl})\\, dM_{\\rm ecl}\\propto M_{\\rm ecl}^{-\\beta}\\,dM_{\\rm ecl}$, where $M_{\\rm ecl}$ is the mass of the embedded star cluster. Moreover, the most massive star in a cluster, $m_{\\rm max}$, is related in a non-trivial way with the cluster mass \\citep{Weidner2013b}; the mass of the most massive star is higher in more massive clusters. Furthermore, observations indicate that higher SFRs lead to the formation of brighter clusters \\citep{Larsen2002}. \\citet{Weidner2004} find that this empirical correlation can be transformed into a relation between the SFR of the galaxy and the maximum embedded cluster mass of the form \\begin{equation} M_{\\rm ecl}^{\\rm max}(SFR)=8.5 \\times 10^4 \\, SFR^{0.75}\\,\\Msun. \\label{eq:MeclSFR} \\end{equation}   The IGIMF is the sum of all the new born stars in all of the star clusters considering the above ingredients , that is, \\begin{equation} \\xi_{\\rm IGIMF}(m,t)=\\int_{M_{\\rm ecl}^{\\rm min}}^{M_{\\rm ecl}^{\\rm max(SFR(t))}}\\xi(m\\leq m_{\\rm max(M_{\\rm ecl})})\\,\\xi_{\\rm ecl}(M_{\\rm ecl})\\,dM_{\\rm ecl}, \\label{eq:IGIMF} \\end{equation} where $\\xi(m\\leq m_{\\rm max(M_{\\rm ecl})})$ is the stellar IMF given by eq. \\ref{eq:4pow}, and $M_{\\rm ecl}^{\\rm min}$ defines the minimum mass of cluster that can be formed in a galaxy, which is a free parameter of the IGIMF model. Now, for this formulation of the IGIMF, it is assumed that star formation in all the star clusters at all epochs is produced with a canonical IMF. However, violent star formation conditions could drive crowding in massive star clusters \\citep{Elmegreen2004, Shadmehri2004}, and the formation of massive stars could be favoured under such conditions, given that the low-mass limit of star formation would be higher. Following preliminary work of \\citet{Marks2012}, \\citet[][WKP11 hereafter]{Weidner2011} consider that, for embedded clusters $M_{\\rm ecl}<2\\times 10^5\\,\\Msun$, the slope of the canonical IMF (eq. \\ref{eq:4pow}) for stars more massive than $1.3\\,\\Msun$ is $\\alpha_3 = 2.35$. For larger cluster masses, the dependence of the stellar IMF slope for stars in this mass range is given by \\begin{equation} \\alpha_{3}(M_{\\rm ecl})=-1.67 \\times {\\rm log}_{\\rm 10}\\left(\\frac{M_{\\rm ecl}}{10^6\\,\\Msun}\\right)+1.05, \\label{eq:alpha3Ecl} \\end{equation} The effect of this assumption in the IGIMF is analysed in WKP11. The high-mass end of the IGIMF is, in general, steeper than the stellar IMF since the formation of low mass stars is favoured. However, if one includes the effect of crowding in massive clusters, as we have seen, the slope of the canonical IMF for massive stars can become more top heavy than the Salpeter IMF for these massive clusters. When a starburst event with high SFR occurs in a galaxy, the formation of massive clusters is favoured and the formation of massive stars is highly enhanced, so the IGIMF can become also more top heavy than the Salpeter IMF for high mass stars .  This formulation leads to the top heavy IGIMF (TH-IGIMF; WKP11). The slope of this TH-IGIMF above $1.3\\,\\Msun$ can then be computed by a least-squares fit to the calculated IGIMF. Hence, the TH-IGIMF can be translated into a power law of a form similar to the canonical stellar IMF used to derive the IGIMF , but with an exponent $\\alpha_{\\rm TH}$ for high mass stars, that is, \\begin{equation} \\xi_{\\rm TH}(m) = k \\left\\{\\begin{array}{ll} \\left(\\frac{m}{m_{\\rm H}} \\right)^{-\\alpha_{1}}&\\hspace{-0.25cm},m_{\\rm H} \\le \\frac{m}{{M}_\\odot} < m_0,\\\\ \\left(\\frac{m_{0}}{m_{\\rm H}} \\right)^{-\\alpha_{1}} \\left(\\frac{m}{m_{0}} \\right)^{-\\alpha_{2}}&\\hspace{-0.25cm},m_0 \\le \\frac{m}{{M}_\\odot} < m_1,\\\\ \\left(\\frac{m_{0}}{m_{\\rm H}} \\right)^{-\\alpha_{1}} \\left(\\frac{m_{1}}{m_{0}} \\right)^{-\\alpha_{2}} \\left(\\frac{m}{m_{1}} \\right)^{-\\alpha_{\\rm TH}}&\\hspace{-0.25cm},m_1 \\le \\frac{m}{{M}_\\odot} \\le m_\\mathrm{max}.\\\\ \\end{array} \\right. \\label{eq:THIGIMF-FIT} \\end{equation} where $m_{\\rm H} = 0.1$, $m_{\\rm 0} = 0.5$ and $m_{\\rm 1} = 1.3$.  After having introduced the basic aspects of the TH-IGIMF theory, we describe the way in which it is implemented in \\sag. Although the reference to ``starbursts condition'' in the work of WKP11 implies a high level of star formation rate in a starburst, we use the resulting dependence of the $\\alpha_{\\rm TH}$ with SFR for any event of star formation in our model, that is, without making any distinction between quiescent star formation mode and starbursts. For each star formation event, we assign to it the TH-IGIMF given by eq. \\ref{eq:THIGIMF-FIT}, according to the corresponding SFR. The model considers small variations of log(SFR) so that the slope of the TH-IGIMF is binned in steps of $0.05$. We adopt $m_{\\rm H}=0.1 M_{\\odot}$ for the lower limit in the TH-IGIMF according to our chemical implementation.   We consider embedded cluster mass functions characterized by different values of the exponent $\\beta$, and of the minimum embedded cluster masses $M_{\\rm ecl}^{\\rm min}$. We explore cases with (i) $\\beta=2$ and $M_{\\rm ecl}^{\\rm min}=5 \\, \\Msun$, (ii)  $\\beta=2.1$ and $M_{\\rm ecl}^{\\rm min}=5 \\, \\Msun$, and (iii) $\\beta=2$ and $M_{\\rm ecl}^{\\rm min}=100 \\, \\Msun$. It is worth noting that some stellar associations with masses as low as $5\\,\\Msun$ are found in the galaxy \\citep[e.g. Taurus-Auriga star-forming regions,][]{Kirk2011} , thus favouring the lowest value of $M_{\\rm ecl}^{\\rm min}$ considered here. Therefore, at least in quiescent star forming conditions, the $M_{\\rm ecl}^{\\rm min}$ is supported observationally. However, this constraint could not be valid in more extreme star-forming conditions.  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{fig1.eps} \\vspace{1pt} \\caption{Dependence of the slope of the TH-IGIMF ($\\alpha_{\\rm TH}$) on the SFR for different minimum embedded cluster masses and slopes of the embedded star cluster mass function: (i) $\\beta=2$, $M_{\\rm ecl}^{\\rm min}=5 \\Msun$ (solid line), (ii) $\\beta=2.1$, $M_{\\rm ecl}^{\\rm min}=5 \\Msun$ (dashed-dotted line), and (iii) $\\beta=2$, $M_{\\rm ecl}^{\\rm min}=100 \\Msun$ (dashed line). } \\label{fig:alphaTHIGIMF} \\end{figure}  Fig. \\ref{fig:alphaTHIGIMF} shows the slope $\\alpha_{\\rm TH}$ as a function of the SFR for different values of $M_{\\rm ecl}^{\\rm min}$ and $\\beta$. It is clear from this plot that the slope $\\alpha_{\\rm TH}$ of the TH-IGIMF is larger than the slope of the Salpeter IMF (characterized by a fixed value of $\\alpha$ for all SFR) in lower star formation regimes. The values of SFR at which the TH-IGIMF becomes more top heavy than Salpeter depend on the value of $\\beta$ and $M_{\\rm ecl}^{\\rm min}$, being of the order of $1,\\,10$ and $1000\\ \\Msun\\,{\\rm yr^{-1}}$ for $\\beta=2$ and $M_{\\rm ecl}^{\\rm min}=100 \\,\\Msun$, $\\beta=2$ and $M_{\\rm ecl}^{\\rm min}=5 \\,\\Msun$, and $\\beta=2.1$ and $M_{\\rm ecl}^{\\rm min}=5 \\,\\Msun$, respectively.  During periods of high star formation activity, the formation of massive stars is favoured. For higher SFRs, the number of high mass stars in a TH-IGIMF becomes larger than in the Salpeter IMF, with the corresponding reduction of the number of low mass stars required by the IMF normalization. This aspect is crucial to understand the build-up of the $[\\alpha/{\\rm Fe}]$-stellar mass relation, as we will see in Subsection \\ref{sec:AlphaFeTHIGIMF}. ",
        "conclusions": ""
    },
    "1402/1402.4964_arXiv.txt": {
        "abstract": "We present a setup that provides a partial UV-completion of the ghost inflation model up to a scale which can be almost as high as the Planck mass. This is achieved by coupling the inflaton to the Lorentz-violating sector described by the Einstein-aether theory or its khronometric version. Compared to previous works on ghost inflation our setup allows to go beyond the study of small perturbations and include the background dynamics in a unified framework. In the specific regime when the expansion of the Universe is dominated by the kinetic energy of the inflaton we find that the model predicts rather high tensor-to-scalar ratio $r \\sim 0.02\\div 0.2$ and non-Gaussianity of equilateral type with $f_{NL}$ in the range from $-50$ to $-5$. ",
        "introduction": "\\label{sec:intro}  All structures in the observed Universe, such as galaxy clusters, galaxies and stars are believed to arise from tiny quantum fluctuations, amplified during a primordial stage in the expansion of the Universe. The most successful model of this stage is the inflationary theory which fits very well all current cosmological data. Generally, inflation is supposed to occur at very high energies (probably a few orders below the Planck mass), where the effects of new unknown physics may show up. Thus the inflationary period provides a unique opportunity to probe these scales through cosmological observations. An intriguing possibility is that the known space-time symmetries become invalid at the inflationary energies. One of such symmetries is Lorentz invariance (LI) which is at the basis of the highly successful Standard Model of particle physics and General Relativity (GR). However, beautiful as it is, GR suffers from the problem of non-renormalizability precluding it from being a consistent quantum theory at energies above the Planck scale. It has been suggested by Ho\\v rava \\cite{Horava:2009uw} that the situation can be improved by allowing LI to be violated at high energy. A consistent implementation of this proposal \\cite{Blas:2009qj} involves additional light degrees of freedom in the gravitational sector that introduce departures from LI even at energies well below Planckian. The description at such energies is provided by the so-called ``khronometric'' model \\cite{Blas:2010hb}, which can be considered as a variant of the phenomenological ``Einstein-aether'' theory \\cite{Jacobson:2000xp,Jacobson:2008aj} for the study of Lorentz-violating (LV) effects in gravity (see \\cite{Jacobson:2010mx,Jacobson:2013xta} for the precise relationship between the two models). The existing astrophysical and cosmological data significantly constrain the parameters of the model \\cite{Elliott:2005va, Blas:2010hb,Blas:2011zd,Audren:2013dwa,Shao:2013wga,Yagi:2013qpa, Yagi:2013ava}, but still leave open a theoretically motivated portion of the parameter space. It is worth mentioning that any application of these ideas to realistic model building must include a mechanism that would prevent significant percolation of LI breaking from gravity into the Standard Model sector where LI has been tested with an outstanding precision; several options have been discussed in \\cite{GrootNibbelink:2004za,Pospelov:2010mp,Pujolas:2011sk,Bednik:2013nxa}.  LV in gravity, though tightly constrained at low energies accessible to current experiments, may be significantly stronger during inflation. An interesting alternative to the standard slow-roll inflation involving LV in gravity is provided by the ghost inflation model \\cite{ArkaniHamed:2003uz} and its ``tilted'' extension \\cite{Senatore:2004rj}. The important feature of these models is the modified dispersion relation for the excitations of the inflaton\\footnote{By inflaton we loosely understand the field generating the primordial perturbations.}. The dispersion relation is quadratic, $\\omega^2\\propto k^4$, in the case of the original ghost inflation and linear, $\\omega^2=\\delta^2\\cdot k^2$, with the small ``sound speed'' $\\delta\\ll 1$ in the tilted version. This, combined with the specific form of the interactions, leads to interesting predictions for the amplitude and shape of non-Gaussianity, which have been constrained by the Planck results \\cite{Ade:2013ydc}.  The ghost inflation models are formulated as effective field theories (EFT) for cosmological perturbations valid below a certain cutoff. While this approach has the advantage of being very general, it is unable to fully capture the evolution of the inflaton background; in particular, it is problematic to incorporate in it the end of inflation and reheating. Besides, if one assumes that the ghost condensate present at inflation persists till today, stability requires the cutoff to be rather low \\cite{ArkaniHamed:2003uy,ArkaniHamed:2005gu}. This makes a UV-completion of ghost inflation desirable.  In this paper we show that ghost inflation can be embedded in the framework of the khronometric or Einstein-aether gravity. The latter being also an EFT, it has a cutoff as well. But this can be as high as just a few orders of magnitude below the Planck mass. In other words, khronometric / Einstein-aether can provide a partial UV-completion (``UV-extension\") for the ghost condensate model almost up to the Planck scale. This allows to describe the evolution of the background and perturbations within a self-contained theory. Interestingly, the UV-extension happens without restoration of the broken Lorentz symmetry, similar to the case recently discussed in \\cite{Endlich:2013vfa}. Further, in the khronometric version of the setup there is a potential UV-completion all the way above the Planck scale in the form of Ho\\v rava gravity.  We study the observational signatures of the extended model. For two reasons the analysis somewhat differs from the discussion of the generic ghost inflation present in the literature. First, the ability to keep the background evolution under control allows to consider the situation when the VEV of the inflaton time derivative $\\langle \\dot\\Theta\\rangle$ --- the ``ghost condensate'' ---  varies with time. This produces a contribution into the tilt of the power spectrum in addition to that coming from a potential for $\\Theta$, which in some cases can actually be dominant. Second, the ghost inflation is characterized essentially by three energy scales: the scale $\\rho_{inf}^{1/4}$ of the inflationary energy density; the scale $M$ of the ghost condensate $M^2=\\langle\\dot\\Theta\\rangle$ which determines the strength of the inflaton self-interaction producing the leading non-Gaussianity; and the scale $M'$ suppressing terms with higher spatial derivatives in the quadratic effective action for inflaton perturbations. In the previous treatments of the ghost inflation the latter two scales have been commonly assumed to be of the same order with the first scale being much higher, $\\rho_{inf}^{1/4}\\gg M\\sim M'$. We will see that in the extended model where all the above scales are derived quantities, $M$ and $M'$ have different parameter dependence, and the requirement that the EFT description is valid for the inflationary background imposes the hierarchy\\footnote{At first sight, this hierarchy might seem surprising from the viewpoint of the EFT for the inflaton {\\it perturbations}. However, it is straightforward to verify that it is stable under radiative corrections and thus perfectly natural.} $M\\gg M'$. On the other hand, $\\rho_{inf}^{1/4}$ can naturally be of the same order as $M$. As a result of this new hierarchy, the amplitude of non-Gaussianity is somewhat suppressed compared to the original prediction of the ghost inflation; still, it remains large enough to be observationally interesting.  The paper is organized as follows. In Sec.~\\ref{sec:fast} we describe the model and identify the inflationary regime that reproduces ghost inflation. In Sec.~\\ref{sec:lin} we study linear cosmological perturbations emphasizing the similarities and differences with the generic ghost inflation treatment. Sec.~\\ref{sec:bispectra} contains calculation of the bispectrum. In Sec.~\\ref{sec:kinetic} we apply our results to the special case when the background dynamics is dominated by the kinetic energy of the inflaton and derive observational constraints on the model parameters in this case. Sec.~\\ref{sec:discussion} is devoted to conclusions. Some details of the analysis are postponed to the Appendices. ",
        "conclusions": "\\label{sec:discussion}  In this paper we presented a setup that provides a partial UV-completion (``UV-extension\") of the well-known ghost inflation model up to a scale which can be almost as high as the Planck mass. This is achieved by coupling the inflaton to the Lorentz-violating sector described by the Einstein-aether theory or its khronometric version. The cutoff of our model coincides with the cutoff of the Einstein-aether sector and is unrelated to the scale of the ``ghost condensate'' --- the dynamically developed expectation value for the time derivative of the inflaton. In the khronometric version the construction can be potentially completed even further in the framework of the Ho\\v rava gravity. Curiously, the UV-extension occurs without restoration of the broken Lorentz symmetry, cf. \\cite{Endlich:2013vfa}.  We have studied the inflationary evolution of the Universe and the generation of primordial perturbations in the model. The latter are described by an effective theory containing a single degree of freedom and governed by the same effective action as in the ghost inflation. The novelty of our model compared to the previous works on ghost inflation is that it allows to go beyond the study of small perturbations and incorporates in a unified framework the background dynamics. We have found that the ghost condensate gives positive contribution into the energy budget of the Universe. This makes inflation possible even in the absence of any potential for the inflaton --- the regime that we called kinetically driven inflation. In principle, it is straightforward to incorporate in our model the graceful exit from inflation and reheating which proceeds through vanishing\\footnote{At least up to the tiny present value of the dark energy density. We do not consider a fine-tuned situation when after inflation the ghost condensate stays large and its positive energy is cancelled by a negative cosmological constant.} of the inflaton potential and the ghost condensate. This would be impossible in the original formulation of the ghost inflation as the low-energy EFT because its cutoff goes to zero in this limit.  On the phenomenological side, we have calculated the characteristics of the power spectrum and bispectrum predicted by the model. Specifically, in the kinetically driven regime the model predicts a rather high tensor-to-scalar ratio and is already constrained by the bounds from the Planck mission. The non-Gaussianity is predicted to be close to the equilateral type with the amplitude ranging from $f_{NL} \\sim -50$ for $r\\sim 0.02$ to  $f_{NL} \\sim -5$ for $r\\sim 0.2$. This is still well within the limits set by Planck. Optimistically, one can expect it to be probed by planned surveys \\cite{Baumann:2008aq,Amendola:2012ys,Andre:2013nfa}.  There are two directions in which our work can be extended. First, additional insight will be gained from the study of higher statistics. In the general EFT formulation of ghost inflation the trispectrum depends on an additional free coupling constant standing in front of the quartic interaction of inflaton perturbations $(\\d_i\\pi)^2 (\\d_j\\pi)^2$ \\cite{Izumi:2010wm}. On the other hand, in our model this coupling is unambiguously determined by the specific form of the UV-extended theory. Repeating the analysis presented in the beginning of Sec.~\\ref{sec:bispectra} it is straightforward to find the leading quartic interaction \\[ S_{[\\pi]}^{(4)}=\\int d^4x\\;\\frac{1}{8\\m^4c_\\T^6 a}\\,(\\d_i\\pi)^2 (\\d_j\\pi)^2\\;. \\] Clearly, the quartic coupling is related the cubic one, see Eq.~(\\ref{Hint}), and thus the shape of the trispectrum will be uniquely predicted.  Second, it is interesting to investigate if the (partial) UV-completion along the lines presented in this paper can be found for other Lorentz-violating models of modified gravity. From the theory viewpoint, the case of Lorentz-violating massive gravity \\cite{Dubovsky:2004sg} deserves particular attention as its completion above the scale of strong coupling associated to the graviton mass could be considered as the analog of the Higgs mechanism in massive Yang--Mills theory."
    },
    "1402/1402.5545_arXiv.txt": {
        "abstract": "We present results of optical broad-band and narrow-band H$\\alpha$ observations of a sample of forty  nearby early-type galaxies. The majority of sample galaxies are known to have dust in various forms viz. dust lanes, nuclear dust and patchy/filamentary dust. A detailed study of dust was  performed for 12 galaxies with prominent dust features. The extinction curves for these galaxies run parallel to the Galactic extinction curve, implying that the properties of dust in these galaxies   are similar to those of the Milky-Way.  The ratio of total to selective extinction  ($R_{V}$) varies between 2.1 to 3.8, with an average of 2.9$\\pm$0.2, fairly close to its canonical value of 3.1 for our Galaxy. The average relative grain size $\\frac{<a>}{a_{Gal}}$ of dust particles in these galaxies turns out to be  1.01$\\pm$0.2, while  dust mass  estimated using optical extinction lies in the range $\\sim$ $10^{2}$ to $10^{4}$ M$_{\\odot}$. The H$\\alpha$ emission  was detected in 23 out of 29 galaxies imaged through narrow-band filters with the H$\\alpha$ luminosities in the range 10$^{38}$ - 10$^{41}$ erg sec$^{-1}$. The mass of the ionized gas is in the range $\\sim$  10$^{3}$ - 10$^{5}$ M$_{\\odot}$. The morphology and extent  of ionized gas  is found similar to those of dust, indicating possible coexistence  of dust and ionized gas in these galaxies. The absence of any apparent correlation between  blue luminosity and normalized IRAS dust mass is suggestive of merger related origin of dust and gas in these galaxies. ",
        "introduction": "Interstellar matter (ISM), being an important factor governing the formation and subsequent evolution of galaxies, has drawn considerable attention since their detection and is still being continued to gain better understanding in various galactic and extragalactic environments. The availability of space observatories/satellites as well as ground-based telescopes with state of the art detectors  has made exploration of different phases of ISM  possible.  From the first detection of dust lanes in early-type galaxies (\\citealt{Ber78}), the possibility of using its orientation  to extract additional information regarding  intrinsic shape of galaxies was explored (\\citealt{Gun79, vanA82, Habe85, Habe88}). Importance of dust in understanding  three dimensional structure of these galaxies led to search  for  dust in several of them (\\citealt{Haw81, Ebn85, Ver88, Sad85}). Subsequently,  deep optical imaging surveys of large sample of early-type galaxies detected dust in a significant fraction of them (\\citealt{Goud94b, Fer99}).   Dust has also been detected in the innermost regions of early-type  galaxies in deep {\\it Hubble Space Telescope} (HST) images, which remained unresolved with the ground-based observations (\\citealt{Jaf94, van95, deKoff00, Tran01}). Further, the detection of significant far-infrared (FIR) emission from early-type galaxies, using {\\it Infra Red Astronomical Satellite} (IRAS), {\\it Infrared Space Observatory} (ISO) and {\\it Spitzer Space Telescope} has added a new dimension in the study of dust and its possible role in underlying dynamics of galaxies (\\citealt{Temi04, Xil04}).  Optical broad-band imaging of galaxies with prominent dust features allows one to investigate physical properties of dust such as  particle size, extinction, reddening,  total dust content of galaxies and the processes that govern their evolution in different environments. Investigation of  dust extinction in different bands {\\it i.e.} extinction curve has been traditionally  a basic tool for studying  dust properties. A refined form of this technique has been applied to study the  properties of dust in a number of dusty early-type galaxies (\\citealt{Goud94a, Sahu98, Pat07, Fin08, Fin10d}). These studies showed that the extinction curves of dust in early-type galaxies run almost parallel to that of the Milky-Way, with average relative grain size not very much different from that in our Galaxy.  Optical spectroscopic observations of early-type galaxies show that $\\sim$55 - 60\\% of them have faint emission line,  indicating the presence of  ionized gas, with mass $\\sim$ 10$^{3}$ -10$^{4}$ M$\\odot$  (Caldwell 1984, Phillips et al. 1986). Further, the H$\\alpha$ imaging survey of large sample of early-type galaxies confirmed the presence of  ionized gas with various morphologies; such as flattened disc, ring or filamentary structures  (\\citealt{Kim89, Shi91, Trin91, Bus93, Goud94a, Sing95, Macch96, Mart04}). The possible sources of gas excitation   have been explored and the post-asymptotic giant branch (pAGB) stars  were identified as the main contributor of ionizing radiation. However,  for at least  10\\% of  early-type galaxies,  the ionized emission is powered by  recently formed stellar subcomponent (\\citealt{Sarzi10}).  Other forms of ISM,  {\\it i.e.} hot and cold  gas have also been found in early-type galaxies. Hot gas halos have been detected  around early-type galaxies with  X-ray observatories  (\\citealt{For85, Can87, Fab89, Osul01, Sar01, Kim03, Kim10}). Cold molecular gas is detected through CO emission in $\\sim$ 22\\% of all early-type galaxies (\\citealt{You11} and references therein).  In a large fraction of dusty early-type galaxies, the morphology and extent of   ionized gas match with that of dust (\\citealt{Goud94b, Fer99, Pat07}) and in some cases with the X-ray emitting region too (\\citealt{Goud98}). This points towards a possible physical connection between hot, warm and cold phases of ISM. The origin and fate of ISM in early-type galaxies is important as it holds clue to the formation and subsequent evolution of early-type galaxies   (\\citealt{Goud95}). The recent studies have shown that  neither internal nor external origin can explain all the observed properties of ISM in these galaxies.  Hence, a  good balance between internal and external origin is suggested as the source of ISM in early-type galaxies.  In the present paper,  we discuss the properties of dust and ionized gas in a sample of forty  low redshift early-type galaxies.  The paper is organized as follows: Observation and data reduction  in Section 2. Section 3 gives the methodology used for analyzing dust and ionized gas, the results have been discussed in Section 4.  We summarize our results in Section 5. \\begin{table} \\begin{scriptsize} \\begin{center} \\caption{Global parameters of sample galaxies} \\begin{tabular}{|lllllcc|} \\hline Object  & RA       & DEC       & Morph.    &$B^{0}_{T}$& $V_{Helio}$& Size \\\\ &(J2000)   &(J2000)    &(RC3)      &           &  (km/s)  & (arcmin)\\\\ \\hline NGC\\,0383& 01:04:39 & \\,36:09:07 & SA0  & 13.38 & 5098 &  1.6x1.4 \\\\ NGC\\,0708& 01:52:46 & \\,36:09:07 & E    & 13.70 & 4855 &  3.0x2.5 \\\\ NGC\\,0720& 01:53:00 & -13:44:19  & E5   & 11.16 & 1745 &  4.7x2.4 \\\\ NGC\\,1052& 02:41:04 & -08:15:21  & E4   & 12.10 & 1510 &  3.2x2.1 \\\\ NGC\\,1167& 03:01:42 & \\,35:12:21 & SA0  & 13.38 & 4945 &  2.8x2.3 \\\\ NGC\\,1199& 03:01:18 & -15:48:29  & E3   & 12.37 & 2570 &  2.4x1.9 \\\\ NGC\\,1395& 03:38:29 & -23:01:40  & E2   & 10.97 & 1717 &  5.9x4.5 \\\\ UGC\\,2783& 03:34:18 & \\,39:21:25 & S0   & 12.99 & 6173 &  1.3x1.2 \\\\ NGC\\,1407& 03:40:11 & -18:34:49  & E0   & 10.70 & 1779 &  4.6x4.3 \\\\ NGC\\,2534& 08:12:54 & \\,55:40:19 & E1   & 13.70 & 3447 &  1.4x1.2 \\\\ NGC\\,2644& 08:41:31 & \\,04:58:49 & S    & 13.31 & 1939 &  2.1x0.8 \\\\ NGC\\,2768& 09:11:37 & \\,60:02:14 & S0   & 10.84 & 1373 &  8.1x4.0 \\\\ NGC\\,2851& 09:20:30 & -16:29:43  & E    & 15    & 5195 &  1.2x0.5 \\\\ NGC\\,2855& 09:21:27 & -11:54:34  & SA   & 12.63 & 1897 &  2.5x2.2 \\\\ NGC\\,3065& 10:01:55 & \\,72:10:13 & SA   & 13.5  & 2000 &  1.7x1.7 \\\\ NGC\\,3115& 10:05:14 & -07:43:07  & S0   & 09.87 & 663  &  7.2x2.0 \\\\ NGC\\,3377& 10:47:42 & \\,13:59:08 & E5   & 11.24 & 665  &  5.2x3.0 \\\\ M\\,105   & 10:47:49 & \\,12:34:54 & E1   & 10.24 & 911  &  5.4x4.0 \\\\ NGC\\,3489& 11:00:18 & \\,13:54:04 & S    & 11.12 & 677  &  3.5x2.0 \\\\ NGC\\,3607& 11:16:54 & \\,18:03:07 & SA   & 10.82 & 960  &  4.9x2.5 \\\\ NGC\\,3801& 11:40:16 & \\,17:43:41 & S0   & 12.96 & 3317 &  3.5x2.1 \\\\ NGC\\,4125& 12:08:06 & \\,65:10:27 & E6   & 10.65 & 1356 &  5.8x3.2 \\\\ NGC\\,4233& 12:17:07 & \\,07:37:28 & S0   & 12.8  & 2371 &  2.3x0.9 \\\\ NGC\\,4278& 12:20:06 & \\,29:16:51 & E1   & 11.20 & 649  &  4.1x3.8 \\\\ NGC\\,4365& 12:24:28 & \\,07:19:03 & E3   & 10.52 & 1243 &  6.9x5.0 \\\\ NGC\\,4494& 12:31:24 & \\,25:46:30 & E1   & 10.71 & 1344 &  4.8x3.0 \\\\ NGC\\,4552& 12:35:39 & \\,12:33:23 & E    & 10.73 & 340  &  5.1x4.7 \\\\ NGC\\,4648& 12:41:44 & \\,74:25:15 & E3   & 12.96 & 1414 &  2.1x1.6 \\\\ NGC\\,4649& 12:43:39 & \\,11:33:09 & E2   & 09.81 & 1117 &  7.4x6.0 \\\\ NGC\\,4697& 12:48:35 &  -05:48:02 & E6   & 10.10 & 1241 &  7.2x4.7 \\\\ NGC\\,4874& 12:59:35 & \\,27:57:34 & cD   & 12.63 & 7224 &  1.9x1.9 \\\\ NGC\\,5322& 13:49:15 & \\,60:11:26 & E3   & 11.14 & 1754 &  5.9x3.9 \\\\ NGC\\,5525& 14:15:39 & \\,14:16:57 & S0   & 13.6  & 5553 &  1.4x0.9 \\\\ NGC\\,5812& 15:00:55 &  -07:27:26 & E0   & 12.19 & 1970 &  2.1x1.9 \\\\ NGC\\,5846& 15:06:29 & \\,01:36:20 & E0   & 11.05 & 1714 &  4.1x3.8 \\\\ NGC\\,5866& 15:06:29 & \\,55:45:48 & S0   & 10.74 & 672  &  4.7x1.9 \\\\ NGC\\,6166& 16:28:38 & \\,39:33:06 & E2   & 12.78 & 9100 &  1.9x1.4 \\\\ NGC\\,7052& 21:18:33 & \\,26:26:48 & E    & 13.40 & 4672 &  2.5x1.4 \\\\ NGC\\,7454& 23:01:07 & \\,16:22:58 & E4   & 12.78 & 2022 &  2.2x1.6 \\\\ IC\\,2476 & 09:27:52 & \\,29:59:09 & S0   & 13.85 & 8007 &  1.5x1.4\\\\ \\hline \\multicolumn{7}{p{12cm}}{ Cols.(2) and (3) list galaxy co-ordinates, Col.(4) lists morphological classification of the galaxies,blue luminosity of the program galaxies are listed in Col.(5), while Col.(6) lists heliocentric velocity, Col.(7) optical size of the galaxies, all taken from RC3(de Vaucoulers et al. 1991)} \\end{tabular} \\label{basic} \\end{center} \\end{scriptsize} \\end{table} ",
        "conclusions": "\\textit{Dust}:   For the galaxies with obvious features of dust extinction, it is found that  morphology and extent of the dust feature in colour-index image is similar to  that of extinction map. This  indicates  that the extinction map  approximates the dust obscured region to a high degree. An analysis of  extinction curves for 12 galaxies is presented  in this work. The extinction curves for eight galaxies are   reported for the first time. The value of $R_{V}$ and relative grain  size for the sample galaxies obtained by analyzing the extinction maps are given in Table 5. The value of  $R_{V}$  for  sample galaxies  varies from 2.1 to 3.8 with an average of 2.95. It is close to the previously reported values by \\cite{Goud94c} ($R_{V}$= 2.7), \\cite{Pat07} ($R_{V}$= 3.01) and \\cite{Fin10d} ($R_{V}$= 2.82). We find that the galaxies having  $R_{V}$ values less than the canonical value of 3.1 have well settled dust morphology in the form of a dust lane or a ring and have relatively smaller  grain  size  than  that of the  Milky Way. Conversely, the galaxies with larger values of  $R_{V}$ show irregular dust morphology and larger  relative grain size. This result is in good agreement with results of \\cite{Goud94b}, \\cite{Pat07} and \\cite{Fin10i}. There are four galaxies in common with the samples of \\cite{Goud94c} and \\cite{Pat07}. Table 8  presents the  value of $R_{V}$, relative grain size and optical dust mass for the four common galaxies  along with the previously reported  values in the literature.   Except for galaxy   \\astrobj{NGC 4125}, which shows complex dust structure (\\citealt{Goud94c}),  our results are in good agreement with those reported in the literature.   \\cite{Goud94b}  explored  various dust destruction mechanisms such as  sputtering of dust by supernova blast wave, turbulent shocks, thermal ions and  hot gas etc. It is  shown that for volatile grain of radius $\\sim$ 0.1 $\\mu$m, the  lifetime of  dust grain varies from 10$^{7}$ to 10$^{9}$ years. The mass of dust and its grain size are expected to drop gradually with time since the dust was acquired by the galaxy.  In this scenario, the smooth and  regular dust lane morphology  associated with smaller grain size is expected, if the galaxy had sufficient time for the dust to settle down in a regular shape.  On the other hand if the dust is acquired recently, it will lack regular dust  morphology as  in the case of \\astrobj{NGC 3801}. Among our sample galaxies \\astrobj{NGC 3801} shows considerably larger relative grain size (1.47) and larger value of $R_{V}$. A multiwavelength study of this galaxy by \\cite{Hota12} revealed that this galaxy has   a kinematically decoupled core or an extremely warped gas disc, dust filaments with complex structures, recent star burst with age less than 500 Myr and  evidence of ionized gas (as seen in this study too). All these facts led \\cite{Hota12} to suggest that this galaxy is a merger-remnant early-type galaxy.  The results of \\cite{Hota12} support that  \\astrobj{NGC 3801}  acquired dust in the recent past and did not have enough time to settle it into a smooth,  regular morphology, which may result in the observed  larger grain size and $R_{V}$ value, as suggested by \\cite{Goud94c}.  The  dust mass reported in this work using optical extinction is in the range $10^2 - 10^4$ $M_{\\odot}$  while, the dust mass estimated using  IRAS fluxes is higher by a factor of $\\sim$ 10$^2$. This indicates  that a significant fraction of dust is diffusely distributed throughout the galaxy (\\citealt{Goud95, Wise96}) which could not be detected with optical observations, this is in agreement with that  of previous studies (\\citealt{Pat07, Fin08, Fin10d}).  \\begin{table*}\\tiny{ \\caption{Comparison of dust properties} \\begin{tabular}{|l|ccl|cccl|} \\hline Galaxy Name&\\multicolumn{3}{|c|}{This Paper}& \\multicolumn{4}{|c|}{Reported values for comparison}\\\\ \\hline & $R_{V}$&$\\frac{<a>}{a_{Gal}}$ & log($\\frac{M_{d,opt}}{M_{\\odot}}$) & $R_{V}$&$\\frac{<a>}{a_{Gal}}$&log($\\frac{M_{d,opt}}{M_{\\odot}}$) & Reference\\\\ \\hline NGC\\,2534&2.1$\\pm$0.20&0.83&3.32$\\pm$0.03&2.03$\\pm$0.28&0.80&3.78&  \\cite{Pat07} \\\\ NGC\\,3489&3.1$\\pm$0.33&1.04&2.40$\\pm$0.05&3.38$\\pm$0.21&1.09&3.66& \\cite{Pat07} \\\\ NGC\\,4125&3.6$\\pm$0.27&1.36&2.10$\\pm$0.02&2.74$\\pm$0.33&0.96&5.22& \\cite{Goud94c}\\\\ NGC\\,5525&3.2$\\pm$0.28&1.01&4.04$\\pm$0.09&3.15$\\pm$0.17&0.99&5.67& \\cite{Pat07} \\\\ \\hline \\end{tabular} \\label{dust_compare} } \\end{table*}   \\textit{Ionized gas}: H$\\alpha$ line emission is detected in $\\sim$ 85\\% of our sample  galaxies observed in narrow-band, the results are listed in Table 7. % The H$\\alpha$ luminosity of the sample galaxies lies in the range 10$^{38}$ - 10$^{41}$ {\\it erg sec$^{-1}$} and the mass of ionized gas lies in the range $\\sim$  10$^{3}$ - 10$^{5}$ M$_{\\odot}$.  The ionized gas and dust have identical morphology  in almost all the  galaxies detected in H$\\alpha$ (refer Figure 2 and 4). % We have several galaxies in common with earlier studies by  \\cite{Kim89}, \\cite{Trin91}, \\cite{Shi91}, \\cite{Goud94b}, \\cite{Macch96} and  \\cite{Fin10i}. The H$\\alpha$ flux estimated in this work are in good agreement with the flux reported in the literature (refer Table 9).% There are 6 galaxies for which H$\\alpha$ flux are available from more than one source and a significant scatter is noticed among them. The source of scatter may be due to the  higher uncertainty involved in the method of determination of the line emission fluxes as discussed in Section 3.3. Similar discrepancy in  H$\\alpha$ flux was also reported by \\cite{Fin10i}.  \\begin{table}\\tiny{ \\begin{center} \\caption{Comparison of ${H\\alpha}$ fluxes} \\begin{tabular}{|llll|} \\hline Galaxy &$\\textit{f}_{H\\alpha}$& Other obs. & Reference \\\\ name& This paper &&\\\\ &$ergcm^{2}s^{-1}$&$ergcm^{2}s^{-1}$&\\\\ \\hline NGC0708& 39$\\times 10^{-14}$  &49$\\times 10^{-14}$& \\cite{Fin10i}\\\\ NGC0720& 15.7$\\times 10^{-14}$&$<$ 1.49$\\times 10^{-14}$& \\cite{Goud94b}\\\\ &&$< 3.53 \\times 10^{-14}$& \\cite{Shi91}\\\\ &&17.66$\\times 10^{-14}$& \\cite{Macch96}\\\\ NGC1052& 22.4$\\times 10^{-14}$&8.1$\\times 10^{-14}$& \\cite{Kim89}\\\\ NGC1199& 3.5$\\times 10^{-14}$ &4.9$\\times 10^{-14}$& \\cite{Fin10i}\\\\ NGC1395& 3.91$\\times 10^{-14}$&2.2$\\times 10^{-14}$& \\cite{Goud94b}\\\\ &&22.7$\\times 10^{-14}$& \\cite{Macch96}\\\\ &&9.33$\\times 10^{-14}$& \\cite{Trin91}\\\\ NGC1407& 18.0$\\times 10^{-14}$&0.9$\\times 10^{-14}$& \\cite{Goud94b}\\\\ &&$<$5.11$\\times 10^{-14}$& \\cite{Shi91}\\\\ &&4.82$\\times 10^{-14}$& \\cite{Macch96}\\\\ NGC2534& 1.13$\\times 10^{-14}$&6.4$\\times 10^{-14}$& \\cite{Fin10i} \\\\ NGC3377& 7.24$\\times 10^{-14}$&11.9$\\times 10^{-14}$& \\cite{Goud94b}\\\\ NGC3489& 15.2$\\times 10^{-14}$&51.89$\\times 10^{-14}$& \\cite{Macch96}\\\\ NGC3607& 35.0$\\times 10^{-14}$&7.83$\\times 10^{-14}$& \\cite{Macch96}\\\\ NGC4125& 14.1$\\times 10^{-14}$&39.5$\\times 10^{-14}$& \\cite{Goud94b}\\\\ &&16$\\times 10^{-14}$& \\cite{Kim89}\\\\ NGC4649& 9.79$\\times 10^{-14}$&11$\\times 10^{-14}$& \\cite{Trin91}\\\\ NGC4697& 31.8$\\times 10^{-14}$&29.5$\\times 10^{-14}$& \\cite{Goud94b}\\\\ &&4.7$\\times 10^{-14}$& \\cite{Trin91}\\\\ NGC5812& 25.3$\\times 10^{-14}$&5.83$\\times 10^{-14}$&\\cite{Macch96}\\\\ NGC5846& 18.9$\\times 10^{-14}$&16$\\times 10^{-14}$& \\cite{Trin91}\\\\ &&28.02$\\times 10^{-14}$& \\cite{Macch96} \\\\ NGC7052& 0.74$\\times 10^{-14}$&6.0$\\times 10^{-14}$& \\cite{Fin10i}\\\\ \\hline \\end{tabular} \\label{h_compare}\\end{center} } \\end{table}   \\subsection{Association of dust and ionized gas} Association of dust and ionized gas and their possible origin in  E/SO galaxies were  studied recently by \\cite{Fin12}. The HST observations revealed presence of nuclear and filamentary/patchy dust in the central part of a  significant fraction of early-type galaxies (\\citealt{Tran01, Mart04}). Most of the galaxies from our sample show dust in their central part.  A study similar to that of \\cite{Fin12} was carried out, by including early-type galaxies  with nuclear, filamentary/patchy dust to the sample of  galaxies having dust lanes. Various relevant parameters like optical dust mass, IRAS dust mass, ionized gas mass, optical luminosity, IR luminosity, H$\\alpha$ luminosity etc. were collected  from \\cite{Bus93}, \\cite{Goud94b}, \\cite{Goud95}, \\cite{Sahu96} \\cite{Sahu98}, \\cite{Dew99}, \\cite{Sing95}, \\cite{Tran01}, \\cite{Pat07}, \\cite{Fin10i}, \\cite{Fin12}, this we denote as combined sample. Wherever necessary the quantities have been normalized for  H$_0$  = 73 km sec$^{-1}$ Mpc$^{-1}$.  The lower optical depth and smaller sampling area in galaxies with patchy /filamentary dust makes it difficult to determine dust mass using extinction values. Dust mass using  optical extinction could be estimated only for 12 galaxies with prominent dust features in our sample. It was demonstrated by \\cite{Tran01}  that even in the case of patchy/filamentary dust, the dust mass determined using  optical extinction and IRAS fluxes are well correlated. Further, the ionized gas mass and  optical dust mass are found to be correlated in early-type galaxies (\\citealt{Fer99, Fin12}). The slope of log(M$_{HII}$) versus  log(M$_{dust, opt}$) for the combined sample was found to be $\\sim$ 0.6, which is shallower than the slope ($\\sim$ 0.8) found by  \\cite{Fin12} for dust lane galaxies.  This indicates that in the prominent dust lane galaxies the source of ionization may be more efficient  compared to the galaxies with patchy/filamentary dust.  In Figure 5 dust mass derived using IRAS fluxes is plotted versus ionized gas mass. A good correlation between these two quantities is seen. The best fitting line to the observed data points with  slope $\\sim$ 0.72 is also shown in the same figure. Our results about the co-existence of dust and ionized gas and their correlation are consistent with those  of similar work by \\cite{Goud94b}, \\cite{Goud95}, \\cite{Macch96} and  \\cite{Fin08, Fin10d}.  \\begin{figure} \\begin{center} \\includegraphics[totalheight=3.4in,width=3.4in]{fig5.eps} \\caption{ Ionized gas mass  plotted versus  dust mass estimated from $IRAS$ fluxes} \\label{mass_ih} \\end{center} \\end{figure}  \\subsection{Possible source of ionization in early-type galaxies} The  common presence of ionized gas in early-type galaxies is well established (\\citealt{Kim89, Goud94b, Macch96, Sarzi06, Fin08}), however, the source of ionization is still debated.  Possible sources of ionization in early-type galaxies were investigated (\\citealt{Mart04, Sarzi10}). \\cite{Sarzi10}  explored various possible sources of ionization e.g. post-asymptotic giant branch  (pAGB) stars, presence of  AGN's, fast shocks, OB stars  and interaction with hot ISM. On the basis of their ionizing balance arguments the pAGB stars were regarded as the best candidate for photoionization; it may be either associated to old stellar population or to recent  star formation. However,  on-going star formation is also responsible for  ionization of gas in at least  $\\sim$ 10\\% of their sample.  The GALEX ultraviolet data revealed that the strong UV flux from  $\\sim$ 30\\%  of nearby bright early-type galaxies can not be explained without invoking $\\sim$ 1 - 3\\% (in total stellar mass) star formation rate in the last billion years. The fraction of galaxies showing recent star formation could be as high as $\\sim$ 50\\%, if corrected for extinction and proper care is taken for  UV flux contribution of AGNs (\\citealt{Kav06}).  Early-type galaxies are known to have  cold molecular gas (\\citealt{Huch95, Mor06, Com07, Oos10}). More recently,  in the ATLAS$^{3D}$ sample of early-type galaxies the detection rate of molecular gas is $\\sim$ 22\\%  with H$_{2}$ mass $\\log$ M(H$_{2}$)/M$\\odot$  in the range  7.10 -  9.29. A strong correlation between  presence of molecular gas and dust, blue features and young stellar ages seen in the H$\\beta$ absorption indicate that the detected molecular gas is often involved in the star formation (\\citealt{You11}). Based on the CO emission in a representative sample of early-type galaxies, \\cite{Com07} and \\cite{Croc11} showed that the  CO-emitting early-type galaxies form a low star formation rate (SFR) extension to the empirical law in spirals.   We searched for available data on CO emission in our sample galaxies. Twelve  out of 26 galaxies with dust and $H{\\alpha}$ emission  in our sample were mapped  to check for CO emission while for remaining  14 galaxies we don't have any information. CO was  detected in only four galaxies namely NGC 2768, NGC 3489 (\\citealt{Croc11})  NGC 3607, NGC 5866 (\\citealt{Davis11}) and upper limits are available for NGC 4125 (\\citealt{wiklind95}), NGC 4649 (\\citealt{Young02}). Using [OIII]/H$\\beta$ ratio \\cite{Croc11} have investigated possible source of ionization in his sample galaxies. It is concluded that in NGC 2768 old pAGB stars/AGN and  in NGC 3489 young pAGB stars are the main source of the ionizing photons. However, to get a better estimate of fraction of CO emitting galaxies, where star formation is the dominant source of ionization, a detailed investigation similar to that done by \\cite{Croc11} is required.  \\subsection{Origin of dust and gas in early-type galaxies} \\begin{figure}\\begin{center} \\includegraphics[totalheight=3.4in,width=3.4in]{fig6.eps} \\caption{Comparison IRAS dust mass, normalized by the luminosity} \\label{mass_lum} \\end{center} \\end{figure} In  early-type galaxies presence of ISM in various forms is well established, but its origin is still an open issue. Many evidences support their co-existence and physical association, pointing  towards a common origin (\\citealt{Goud94b, Goud95, Macch96, Sarzi06, Fin10i, Fin12}). As discussed earlier, $H{\\alpha}$ emission is detected in a significant fraction of our sample galaxies. The matching morphology of line emitting regions with that of optical extinction indicates co-existence of dust and ionized  gas in them. The possible sources for origin of dust and gas in these systems are internal and external. In  case of internal origin dust is deposited through  mass loss from  evolved stars. The dust particles  also get destroyed due to  sputtering in various environments. It  has been demonstrated that the mass of dust observed in early-type galaxies is generally higher than that expected from stellar mass loss, indicating  an external origin of dust such as galaxy interaction or mergers (\\citealt{Goud95, Pat07, Fin12} and references therein).   To investigate possible  origin of dust in early-type galaxies relationship between optical luminosity and  dust mass normalized by  luminosity is  also explored (\\citealt{Goud95, Fer99, Pat07}).  In Figure 6, % IRAS dust mass normalized with respect to blue luminosity  is plotted against  blue luminosity for the combined sample. The absence of any apparent correlation between these quantities suggests the possibility of external  such as merger related origin of dust and gas in these galaxies.  The observed misalignment between the angular momentum vectors of interstellar gas and stellar system in several early-type galaxies suggests that they are kinematically decoupled (\\citealt{Ber88, Kim89, van95, Caon00, Kraj08, Sarzi06}). Further, \\cite{Sarzi06} showed that the kinematical misalignment between the gaseous and stellar component is  strongly dependent on apparent flattening and the level of rotational support in  these galaxies. The flatter and faster rotating early-type galaxies are known to have preferentially co-rotating gaseous and stellar systems, indicating that  an internal origin of gas may play an important role in fast rotating galaxies. Hence, it is necessary  to invoke a balance between the internal and external origin of gas and dust to  explain the recent  observational results (\\citealt{Fin12}). Moreover, the observed blue  colours  (near-UV - $r$ $<$ 5.5) in a large sample of early-type galaxies can only be explained by introducing some amount of recent star formation in these galaxies (\\citealt{Kav06}). The  recycled gas from stellar mass loss within the galaxy is not enough to account for their observed   blue colour.  Hence,  an additional source of  fuel for star formation is unavoidable, may be from  external origin."
    },
    "1402/1402.2686_arXiv.txt": {
        "abstract": "We present IRAM Plateau de Bure Interferometer observations of the CO(3--2) and CO(5--4) line transitions from a \\lya blob at $z$ $\\sim$ 2.7 in order to investigate the gas kinematics, determine the location of the dominant energy source, and study the physical conditions of the molecular gas. CO line and dust continuum emission are detected at the location of a strong MIPS source that is offset by $\\sim$1.5\\arcsec\\ from the \\lya peak. Neither of these emission components is resolved with the 1.7\\arcsec\\ beam, showing that the gas and dust are confined to within $\\sim$7\\,kpc from this galaxy.  No millimeter source is found at the location of the \\lya peak, ruling out a central compact source of star formation as the power source for the \\lya emission. Combined with a spatially-resolved spectrum of \\lya and \\heii, we constrain the kinematics of the extended gas using the CO emission as a tracer of the systemic redshift.  Near the MIPS source, the \\lya profile is symmetric and its line center agrees with that of CO line, implying that there are no significant bulk flows and that the photo-ionization from the MIPS source might be the dominant source of the \\lya emission.  In the region near the \\lya peak, the gas is slowly receding ($\\sim$100\\kms) with respect to the MIPS source, thus making the hyper-/superwind hypothesis unlikely. We find a sub-thermal line ratio between two CO transitions, \\ICO{5}{4}/\\ICO{3}{2} = 0.97 $\\pm$ 0.21. This line ratio is lower than the average values found in high-$z$ SMGs and QSOs, but consistent with the value found in the Galactic center, suggesting that there is a large reservoir of low-density molecular gas that is spread over the MIPS source and its vicinity. ",
        "introduction": "\\setcounter{footnote}{0}  Giant Ly$\\alpha$ nebulae (also known as ``\\lya blobs'') are large (50--100~kpc) spatially extended regions emitting copious amounts of Ly$\\alpha$ emission [$L$(\\lya)$\\sim$ 10$^{43-44}$ \\unitcgslum] \\cite[e.g.,][]{Keel99, Steidel00, Francis01, Matsuda04, Matsuda11, Dey05, Saito06, Smith&Jarvis07, Ouchi09, Prescott09, Prescott12a, Yang09, Yang10, Erb11}.  They may represent sites of massive galaxy formation and their early interaction with the intergalactic medium. However, the questions of what powers these gigantic gas halos and whether the surrounding gas is outflowing from or infalling into the embedded galaxies are still debated. The proposed scenarios include photo-ionization by AGNs \\citep{Geach09}, shock-heated gas by galactic superwinds \\citep{Taniguchi&Shioya00}, cooling radiation from cold-mode accretion \\citep{Fardal01, Haiman00, Dijkstra&Loeb09, Goerdt10}, and resonant scattering of \\lya from star-forming galaxies \\citep{Steidel11, Hayes11}. However, observations of \\lya blobs currently allow no firm conclusions about their nature. For example, less than 20\\% of blobs are found to contain X-ray luminous AGN \\citep{Geach09} implying that a powerful AGN is not always required for producing the extended halo \\cite[see also][]{Yang09}.  Yet \\citet{Overzier13} argue strongly that most are powered by AGN.  It has long been believed that \\lya blobs undergo intense dusty starbursts like submillimeter galaxies \\citep{Ivison98, Chapman04, Geach05}, but recent studies show that most \\lya blobs are not as luminous at rest-frame far-infrared (FIR) wavelengths as submillimeter galaxies \\citep{Yang12, Tamura13}.   The molecular gas content in \\lya blobs has to date remained unconstrained, despite numerous attempts to detect CO emission \\citep{Chapman04, Yang12, Wagg12a}.  Observations of the molecular gas component in these systems are key to determining the physical conditions in the star-forming gas, the star formation efficiency, and constraining various relevant timescales and the role of feedback. CO emission lines are also a useful tool as a tracer of the systemic redshift, without which it is difficult to constrain the gas kinematics within a given system.  For example, previous studies of the same \\lya blob using only \\lya emission line have led to radically different conclusions about the gas kinematics, suggesting that gas is outflowing in stellar or AGN winds \\citep{Wilman05}, inflowing due to gas accretion \\citep{Dijkstra06b}, or even static \\citep{Verhamme08}. These discrepancies mainly arise because the systemic velocities of the galaxies within the \\lya blob are not directly measured, but rather need to be assumed \\cite[see also][]{Bower04,Weijmans09}. While other nebular lines in the UV/optical such as \\oiii and \\ha have been used to constrain the systemic velocity within a small sample of \\lya blobs \\cite[e.g.,][]{Yang11,McLinden13}, they are strongly susceptible to extinction and are only detectable from the ground over a restricted redshift range due to atmosphere.  In such cases, the far infrared emission lines such as CO and \\cii 158\\um are ideal probes to determine the systemic redshift of the system, should they harbor a cold molecular component.   In this paper, we discuss new Plateau de Bure Interferometer (PdBI) observations of the CO $J$ = 3\\,$\\rightarrow$\\,2 and $J$ = 5\\,$\\rightarrow$\\,4 line transitions from one of the most luminous \\lya blobs \\cite[LABd05;][]{Dey05}. Originally discovered due to its strong {\\it Spitzer} MIPS 24\\um flux, LABd05 has been the subject of many multi-wavelength studies. {\\it Spitzer} and high resolution {\\sl HST} optical/NIR observations show that this giant \\lya nebula is composed of a strong MIPS source and $\\sim$17 small compact galaxies \\citep{Prescott12b}.  The peak of the \\lya emission is substantially offset from all the galaxies (e.g., 1.5\\arcsec\\ from the MIPS source), but is coincident with a detection in \\heii emission.  \\citet{Prescott08} found that LABd05 inhabits an overdense environment, suggesting that it may be the progenitor of a rich galaxy group or low-mass galaxy cluster. The SED of the bright MIPS source within LABd05 can be explained by an AGN-dominated template with a FIR luminosity of \\lfir(40--1000\\um) = 4\\E{12}\\lsun \\citep{Dey05, Bussmann09, Yang12}, although it is not yet clear whether the far-IR luminosity is driven primarily by star-formation or AGN activity \\citep{Colbert11}. The MIPS source contributes the bulk of the bolometric luminosity from the region and has a very extreme rest-frame UV-to-mid-IR color that characterizes it as a heavily dust-obscured galaxy \\cite[DOG;][]{Dey08, Prescott12b}. On the other hand, the kinematics of \\lya-emitting gas have not been fully constrained.  While the Keck Low Resolution Imaging Spectrometer (LRIS) longslit spectroscopy shows that there is a monotonic (and approximately linear) velocity gradient across the \\lya blob (maybe due to rotation, outflow or infall), the interpretation of this signature was ambiguous because of the lack of the systemic redshift of the embedded galaxies, possibly the center of the gravitational potential of the blob \\citep{Dey05}. Previous attempts to detect the CO emission in LABd05 using the IRAM-30m yielded only upper limits \\citet{Yang12}.  Here, we present new PdBI interferometric observations of LABd05 that reveal detections of the dust continuum emission and CO emission lines from the system.  This paper is organized as follows.  In \\S\\ref{sec:observation}, we describe our PdBI observations.  In \\S\\ref{sec:detection}, we present the detection of molecular gas from LABd05.  In \\S\\ref{sec:voffset}, we constrain the gas kinematics of the \\lya-emitting gas using the systemic redshift derived from CO and \\heii lines.  In \\S\\ref{sec:COSED}, we use the CO line SED to study physical conditions in star-forming regions. In \\S\\ref{sec:SED_LABd05}, we present the FIR SED and derive constraints on the dust properties of LABd05.  In \\S\\ref{sec:model}, we discuss some plausible physical models for LABd05.  Section \\ref{sec:conclusion} summarizes our conclusions. Throughout this paper, we adopt the following cosmological parameters: $H_0$ = 70\\,${\\rm km\\,s^{-1}\\ Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:conclusion}  We have obtained IRAM PdBI observations of the CO $J$ = 3\\,$\\rightarrow$\\,2 and $J$ = 5\\,$\\rightarrow$\\,4 line transitions from a $z$=2.7 \\lya blob (LABd05) in order to investigate the molecular gas content and kinematics, to determine the location of the dominant energy source, and to study the physical conditions of star-forming regions within the \\lya blob.  We detect CO line emission from the molecular gas associated with this \\lya blob. % The CO line emission and the dust continuum are detected at the location of a strong MIPS source, which is offset by $\\sim$1.5\\arcsec\\ (12\\,kpc in projection) from the peak of \\lya surface brightness distribution. Neither the CO line nor the dust continuum emission is resolved with our 1.7\\arcsec\\ beam, showing that the molecular gas and dust are confined to within a $\\sim$7kpc region around the MIPS source. In addition, no millimeter continuum source is found at the location of the \\lya peak, excluding the presence of a compact source of hidden star formation at SFR $>$ 260\\msun{yr$^{-1}$} which might be directly responsible for the \\lya emission.  The CO line spectra show that the MIPS source is indeed located at the same redshift as the extended \\lya emission and that it is a massive galaxy ($M_{\\rm dyn}$ = 3--7\\E{11}\\msun) based on the broad CO line-width (FWHM = 700\\,\\kms).  Combined with Keck/LRIS longslit spectroscopy of \\lya and \\heii, we constrain the kinematics of the extended gas using the CO emission as the best tracer of the systemic redshift.  At the position of the MIPS source, the \\lya profile is symmetric and its line center agrees with those of the CO lines.  This implies that there are no significant bulk flows and the photo-ionization from the MIPS source might be the dominant source of the \\lya emission.  Near the peak of the \\lya nebula, the gas is slowly receding ($\\sim$100\\kms) with respect to the MIPS source, thus disfavoring the hyper-/superwind hypothesis where extreme galactic winds are responsible for the extended \\lya emission.  However, we note that we cannot rule out the possibility that gas flow within LABd05 is tangential to our LOS.  We find that a significantly sub-thermal line ratio between the two CO transitions, \\ICO{5}{4}/\\ICO{3}{2} = 0.97 $\\pm$ 0.21 [\\LpCO{5}{4}/\\LpCO{3}{2} = 0.35 $\\pm$ 0.08]. This line ratio is lower than the average values found in high-$z$ SMGs and QSOs, but is consistent with the value found in the center of the Milky Way.  This line ratio indicates that there is a large reservoir of low-density molecular gas that could be spread over the vicinity of the MIPS source.  Observations of CO(7--6) and CO(1--0) lines with higher spatial resolution are required to further constrain the properties of the star-forming regions within this \\lya blob.    \\medskip"
    },
    "1402/1402.7295_arXiv.txt": {
        "abstract": "To illustrate the complementarity of the linear collider and astrophysics bounds on the light (MeV-scale mass) dark matter (DM), we study the constraints on the magnetic dipole DM from the DM-electron interactions at the proposed International Linear Collider (ILC) and in supernova (SN) 1987A. We in particular focus on the $e^+ e^-$ annihilation which is the common process for producing DM pairs both at the ILC and in the SN. We estimate the bounds on the DM magnetic dipole moment from the mono-photon signals at the ILC and also from the energy loss rate due to the freely streaming DM produced in the SN. The SN bounds can be more stringent than those from the ILC by as much as a factor ${\\cal O}(10^5)$ for a DM mass below $10^2$ MeV. For larger DM masses, on the other hand, SN rapidly loses its sensitivity and the collider constraints can complement the SN constraints. ",
        "introduction": "The nature of the dark matter (DM) remains an outstanding question which can provide crucial clues for the physics beyond the Standard Model (SM). In particular, besides the commonly studied weakly interacting massive particles (WIMPs) with weak scale mass, there has been growing interest in light DM whose parameter region has not yet been experimentally fully explored. For instance, current direct DM search experiments have recoil energy sensitivity down to of order a keV which limits the DM mass to be larger than about a GeV, and it would be of great interest to investigate the lighter mass range below a GeV for the potential window to new physics beyond the SM. Facing the wide open possibilities for the properties of DM, we study the interaction of MeV scale DM particles which possess a magnetic dipole moment and therefore interact with the photon. The DM magnetic dipole moment can be easily generated in many extensions of the SM such as asymmetric DM models and there have been many studies of the dipole DM, in particular for the light DM whose interactions with the SM particles can enjoy infrared enhancement due to the small momentum transfer in the photon exchange \\cite{pos,kam,mas,kopp,dede,bank,zure2,ess,semi,light,barg2,heo3,paolo3,essi,ilidio,gary,nob}.  We aim to illustrate the complementarity of linear collider and astrophysical probes on light dark matter in the MeV-scale mass range. We focus on the interactions of the DM and electron/positron pairs and estimate how they can can affect the ILC and SN signals. For the ILC, we study the impact of the magnetic dipole DM on the mono-photon events where pairs of DM particles arise from $e^+ e^-$ annihilations, and, for the SN, we calculate the DM emission rate potentially affecting  SN cooling which is also due to  pairs of DM particles produced by $e^+ e^-$ annihilations. We relate the collider and SN phenomenology through this common DM production channel of $e^+ e^-$ annihilations and clarify how collider and SN signals can complement each other in constraining the DM magnetic dipole moment. ",
        "conclusions": ""
    },
    "1402/1402.0010_arXiv.txt": {
        "abstract": "Hybrid-kinetic numerical simulations of firehose and mirror instabilities in a collisionless plasma are performed in which pressure anisotropy is driven as the magnetic field is changed by a persistent linear shear $S$. For a decreasing field, it is found that mostly oblique firehose fluctuations grow at ion Larmor scales and saturate with energies $\\propto$$S^{1/2}$; the pressure anisotropy is pinned at the stability threshold by  particle scattering off microscale fluctuations. In contrast, nonlinear mirror fluctuations are large compared to the ion Larmor scale and grow secularly in time; marginality is maintained by an increasing population of resonant particles trapped in magnetic mirrors. After one shear time, saturated order-unity magnetic mirrors are formed and particles scatter off their sharp edges. Both instabilities drive sub-ion-Larmor--scale fluctuations, which appear to be kinetic-Alfv\\'{e}n-wave turbulence. Our results impact theories of momentum and heat transport in astrophysical and space plasmas, in which the stretching of a magnetic field by shear is a generic process. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0592.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {The atmospheric composition of transiting exoplanets can be characterized during transit by spectroscopy. Detections of several chemical species have previously been reported in the atmosphere of gaseous giant exoplanets. For the transit of an Earth twin, models predict that biogenic oxygen (O$_2$) and ozone (O$_3$) atmospheric gases should be detectable, as well as water vapour (H$_2$O), a molecule linked to habitability as we know it on Earth.} % aims heading (mandatory) {The aim is to measure the Earth radius versus wavelength $\\lambda$ - or the atmosphere thickness $h(\\lambda)$ - at the highest spectral resolution available to fully characterize the signature of Earth seen as a transiting  exoplanet.} % methods heading (mandatory) {We present observations of the Moon eclipse of December 21, 2010. Seen from the Moon, the Earth eclipses the Sun and opens access to the Earth atmosphere transmission spectrum. We used two different ESO spectrographs (HARPS and UVES) to take penumbra and umbra high-resolution spectra from $ \\approx 3100$ to 10400\\AA. A change of the quantity of  water vapour above the telescope  compromised the quality of the UVES data. We corrected for this effect in the data processing. We analyzed the data by three different methods. The first method is based on the analysis of pairs of penumbra spectra. The second makes use of a single penumbra spectrum, and the third of all penumbra and umbra spectra.  } % results heading (mandatory) {Profiles $h(\\lambda)$ are obtained with the three methods for both instruments. The first method gives the best result, in agreement with a model. The second method seems to be more sensitive to the Doppler shift of solar spectral lines with respect to the telluric lines. The third method makes use of umbra spectra which bias the result by increasing the overall negative slope of $h(\\lambda)$. It can be corrected for this \\textit{a posteriori} from results with the first method. The three methods clearly show the spectral signature of the Rayleigh scattering in the Earth atmosphere and the bands of H$_2$O, O$_2$, and  O$_3$. Sodium is detected. Assuming no atmospheric perturbations, we show that the E-ELT is theoretically able to detect the $O_2$ A-band in 8~h of integration for an Earth twin at 10~pc.} % conclusions heading (optional), leave it empty if necessary {Biogenic $O_2$, $O_3$, and water vapour are detected in Earth observed as a transiting planet, and, in principle, would be within reach of the E-ELT for an Earth twin at 10~pc.}   % ",
        "introduction": "Observing planetary transits gives access to a set of properties of the  star+planet system (e.g. \\cite{winn2009}). The relative size of the planet with respect to the star, given by the depth of the transit light-curve, is one of the most straightforward methods. Spectrally resolved observations have shown for gaseous giant exoplanets that transit the brightest stars that the transit depth varies with wavelength (the first to report this were \\cite{charbonneau2002},  \\cite{vidal-madjar2003},  \\cite{vidal-madjar2004}) which reveal the signature of the planetary atmosphere. At certain wavelengths, the atmosphere indeed absorbs the stellar light and becomes opaque, which makes the planet look larger in front of the star and produces a deeper transit. The atmosphere of smaller terrestrial exoplanets close to or in the habitable zone (HZ) is in principle also accessible by this method (early works by \\cite{schneider1992}, \\cite{schneider1994}, and more recent models by \\cite{ehrenreich2006}, \\cite{kaltenegger2009}, \\cite{ehrenreich2012}, \\cite{garciamunoz2012}, \\cite{snellen2013} and \\cite{betremieux2013}). Recent simulations show that the next generation of instruments in the infrared or visible domains, that is the James Webb Space Telescope (JWST) and the European Extremely Large Telescope (E-ELT), should be able to offer a first access to characterizing terrestrial planets in the HZ and detecting biosignatures (here atmospheric O$_2$ and O$_3$) and molecules linked to habitability (\\cite{rauer2011}; \\cite{hedelt2013}; \\cite{snellen2013}).  In this context, it is of prime interest to observe Earth as a transiting exoplanet.  Vidal-Madjar et al. (2010) observed the August 2008 partial lunar eclipse with the high-resolution SOPHIE \\'echelle spectrograph at the 193~cm Haute-Provence telescope (\\cite{perruchot2008}; \\cite{bouchy2009}) and showed that the Earth atmosphere thickness versus wavelength can be derived from such observations. Earth observed from the Moon during a lunar eclipse indeed transits in front of the Sun and gives the possibility to study the Earth atmosphere as during a transit. These first observations also allowed one to identify molecular species in the atmosphere, that is water vapour (H$_2$O), a molecule linked to Earth habitabilty, biogenic oxygen (O$_2$), and ozone (O$_3$).  Ozone and oxygen are considered as reliable biosignatures for Earth-like planets (\\cite{owen1980}; \\cite{leger1993}; \\cite{seager2013}) except for terrestrial planets outside the liquid-water habitable zone (\\cite{schindler2000}). These species were also detected by Pall\\'e et al. (2009) from umbra spectra taken during the same eclipse. However, Vidal-Madjar et al. (2010) pointed out that the Pall\\'e et al. spectrum is not the transmission spectrum of Earth observed as a distant transiting planet. The umbra signal indeed results from photons that are highly scattered and refracted by the deep atmosphere. These photons strongly deviate from the star-planet direction and cannot be collected during the observation of a transit.  From penumbra observations, Vidal-Madjar et al. (2010) additionally detected   the thin high-altitude sodium Na I layer. The larger diameter of Earth in the blue due to increasing Rayleigh scattering at shorter wavelengths is also well visible, although it is partially blended with residuals from spectrograph orders correction.  In this paper, we analyze new observations obtained during the lunar eclipse on Dec. 21, 2010 with the ESO spectrographs HARPS (from $ \\approx 3800$ to 6900\\AA, \\cite{mayor2003}) and UVES (from $\\approx 3200$ to 10400\\AA, \\cite{dekker2000}).  We present first the principle of the observation and different reduction methods:  the original Vidal-Madjar et al. (2010) data reduction method (method~1), which makes use of a pair of penumbra spectra, and a second method (method~2), which requires only one single penumbra spectrum to recover the atmosphere thickness profile (Sect.~\\ref{principle}). We also present a third alternative method to derive the thickness profile, which makes use of penumbra and umbra spectra. This third approach is biased by the refraction in the Earth atmosphere (Sect.~\\ref{method3}, method~3), but can be corrected \\textit{a posteriori} with the result of method~1. An observation log and more details about the spectrographs are given in Sect.~\\ref{log}. We give details on the data reduction in Sect.~\\ref{HARPS-data-processing} and Sect.~\\ref{UVES-data-processing} for HARPS and UVES, respectively.  The altitude profiles versus wavelength obtained with the three different methods are detailed in Sect.~\\ref{Results}, as well as simulations of the E-ELT observing a transiting Earth twin at 10~pc.    %__________________________________________________________________ ",
        "conclusions": "We have described the analysis of the observations of a lunar eclipse used as  a proxy to observe Earth in transit. The results fully confirm those obtained with SOPHIE and the observation of the August 2008 eclipse  (\\cite{vidal-madjar2010}) and extend the observed $h(\\lambda)$ profile towards the near infrared thanks to the UVES data. They agree well with the models (\\cite{ehrenreich2006}; \\cite{kaltenegger2009}; \\cite{snellen2013}; \\cite{betremieux2013}). The Rayleigh increase in the blue is visible. Biogenic oxygen and ozone are also well visible in the atmosphere altitude profiles, and are therefore relevant detectable biosignatures to look for during transits of Earth-like exoplanets in the [4000-10500]\\AA\\ range. The oxygen A and B bands show up clearly with HARPS or UVES.  Water vapour is well detected at $\\lambda>$~8000\\AA\\  with UVES because water-vapour absorption is stronger at these wavelengths. Therefore even low densities of water vapour at these altitudes are detected. Water vapour is seen at an altitude of up to 20~km around 9500\\AA. Moreover, Rayleigh scattering and ozone absorption are essentially absent  above 8000\\AA\\, therefore we are also able to probe deeper layers in the Earth atmosphere where low-altitude water vapour is also present at higher densities.   In HARPS data at $\\lambda<$~7000\\AA\\, water vapour appears above 25~km with method~1 but remains below 25~km with method~2 and below 12~km with method~3, because water-vapour absorption is weaker in the visible range.  The results from UVES are nevertheless less reliable than those of HARPS, because the PWV was not stable during the beginning of the night when the calibration spectra were recorded on the full-Moon before the eclipse, which compromised the quality of the eclipse observation. Although we were able to recover the water-vapour content in the calibration spectra \\textit{a posteriori}, the results for the detection of water vapour in the altitude profiles with UVES are not coherent between method~1, 2 and 3. This underlines the difficulties of these observations which were made through the local atmosphere above the telescope, while our goal was to measure atmospheric features at a planetary scale. Lunar eclipse observations done from space, either at low or high spectral resolution, would not suffer from the telluric pollution, with the additional benefit of an access to UV and the prominent Hartley band of ozone.  Another bias in the recovering of the Earth atmosphere profile from a lunar eclipse is the effect of refraction in method~3 that makes use of refracted light from the umbra. We have shown that this effect can be partially corrected with the results from method~1, but the profiles - either from HARPS or UVES - still show a residual negative slope with respect to the profile from method~1 or 2.  Finally these results provide input for the preparation of future observations with the next generation of giant telescopes. We showed that the E-ELT observing in the visible range at low spectral resolution and in the absence of atmospheric perturbations is in principle able to detect the O$_2$ A-band in the atmosphere of a transiting Earth twin. It will be extremely challenging in practice through the atmosphere. If we overcome this difficulty, especially with high spectral resolution spectroscopy, the O$_2$ signature may then be increasing from year to year (from transit to transit), as will  the enthusiasm of the observers."
    },
    "1402/1402.5635_arXiv.txt": {
        "abstract": "We describe the data flow in the operation of the VEGA/CHARA instrument. After a brief summary of the main characteristics and scientific objectives of the VEGA instrument, we explain the standard procedure from the scientific idea up to the execution of the observation. Then, we describe the different steps done after the observation, from the raw data to the archives and the final products. Many tools are used and we show how the Virtual Observatory principles have been implemented for the interoperability of these softwares and data bases. ",
        "introduction": "\\noindent Optical long baseline interferometry suffers from an image of a complex observing technique reserved to specialists of instrumentation. With the advent of widely open large facilities such as VLTI and KECK, this has change a lot these recent years. Interferometric instruments are now not only used by small teams but are widely shared in open time competition. This is true also on the Center for High Angular Resolution Astronomy (CHARA) Array \\citep{chara} where our instrument, VEGA (Visible spEctroGraph and polArimeter) \\citep{mourard09,Mourard2011}, has been in operation for many years. This effort of opening the access to a wider community has driven some of the developments we have done these last years.  An interferometric program relies on an interferometric infrastructure and on one or many instruments. VEGA operates either with $2$, $3$ or $4$ telescopes depending on the science objectives. In all these modes, the choice of the telescopes among the $6$ that CHARA is operating opens a large number of possibilities. It's clear that, starting with a scientific idea, the problem of defining the best observing strategy is not obvious and needs tools. This is usually done by community softwares that simulate the interferometer. Moreover, doing the observations and optimizing the night with the different programs, is also a complicated task. In order to facilitate these two main aspects of the VEGA operation, we have developed a global vision of the information and data flow, starting from the scientific idea up to the final data.  This paper aims to present the different steps of the VEGA use. We present the VEGA instrument and its main science programs in Section~\\ref{vega}. Then Section~\\ref{prepa} presents the preparation phase, Section~\\ref{observation} details the way the observations are conducted and finally Section~\\ref{processing} presents the necessary steps from the raw data to the final products. These steps are schematically summarized in Fig.~\\ref{fig:PIVOT} where the relations between all the software tools and the databases are presented.  \\begin{figure} \\centering \\includegraphics[scale=0.4]{fig1.jpg} \\caption{Diagram of the relations between all the tools used from the preparation of observations to the exploitation of data. The arrows represent the data flow between the different applications.}\\label{fig:PIVOT} \\end{figure} ",
        "conclusions": "In this paper we have described the sequence of actions one has to do to proceed from an interesting scientific idea using high angular resolution up to the resolution of the question with data and comparison with models. Due to the intrinsic complexity of the interferometric measurements, mainly due to the possible large number of observing modes, it has been of utmost importance, for non specialists, to highly simplify the procedure. This was exactly the aim of the work we have developed in the last years and the fact that, now, non-interferometric astronomers are already using VEGA for their scientific programs is certainly a great step forward. Thanks to the development of interoperability by the JMMC and OCA teams, in the framework of the development of the Virtual Observatory principles, it has been possible to easily link the different applications and to use the powerful capabilities of databases to smooth all the interferometric process."
    },
    "1402/1402.4891_arXiv.txt": {
        "abstract": "Different pieces of observational evidence suggest the existence of disks around isolated neutron stars. Such disks could be formed from supernova fallback when neutron stars are born in core-collapse supernova explosions. Efforts have been made to search for disks around different classes of pulsars, which include millisecond pulsars, young neutron star classes (magnetars, central compact objects, and X-ray dim isolated neutron stars), and regular radio pulsars. We review the main results from observations at wavelengths of from optical to sub-millimeter/millimeter. ",
        "introduction": "\\label{sec:intr}  It has long been suggested that isolated pulsars (PSRs) might have disks and the existence of disks would affect evolution of pulsars and help explain some observed features in pulsars \\citep{md81}. The discovery of a planetary system around PSR B1257+12 \\citep{wf92}, which actually was the first discovered extrasolar planetary system, motivated numerous studies about how such a planetary system could be formed. The proposed formation mechanisms more or less involved the existence of an accretion/debris disk around the neutron star (see \\citealt{mh01} and references therein). To date, two additional pulsars, B1620$-$26 and J1719$-$1438, have been found with planetary systems, although the first is a triple system in the globular cluster M4 likely having been formed from a dynamical exchange interaction (\\citealt{sig+03} and references therein), and in the second the planet-mass companion is probably the leftover of a carbon white dwarf having lost most its mass during the phase of low-mass X-ray binary evolution \\citep{bai+11}. In addition to these three pulsars, \\citet{sha+13} recently have shown that the observed, long-term timing variations for PSR B1937+21 are possibly explained by considering an asteroid belt around the pulsar.  The above neutron stars are old (ages $> 10^8$ yrs), so-called recycled pulsars with fast, millisecond spin periods. For regular pulsars formed from core-collapse supernovae for 10$^3$--10$^6$ yrs, rotational spin noise (i.e., timing noise) in them is several orders of magnitude larger than that in millisecond pulsars (MSPs) (e.g., \\citealt{sc10}). It is thus difficult to detect planetary bodies around them from pulsar timing.  However certain phenomena have been seen hinting the existence of accretion/debris disks. \\citet{cs08} summarize four types of pulsar pulse emission variations: nulling, transient pulse emitting from so-called rotating radio transients (RRATs; \\citealt{mcl+06}), subpulse drifting, and emission mode changing, and consider that these phenomena are caused by migration of circumpulsar debris material into the magnetospheres of pulsars. In individual pulsars, for example, the recent measurement of the second period derivative of the spin of PSR J1734$-$3333 suggests that the magnetic field of this pulsar is increasing \\citep{esp+11}, or alternatively its spin properties might be affected by having an accretion disk \\citep{cal+13}. It has also been suggested that because jets are generally associated with accretion disks, young pulsars seen with jets might harbor accretion disks \\citep{bp04}.  While it is not very clear how disks around MSPs would be formed (e.g., \\citealt{sha+13}), particularly for the planetary system case around B1257+12 \\citep{mh01}, young pulsars are believed to possibly have disks due to supernova fallback \\citep{che89}. During a core-collapse supernova explosion, part of ejected material may fallback and if the material has sufficient angular momentum, a disk might be formed around the newly born neutron star \\citep*{lwb91}. The existence of fallback disks has been suggested to be the cause of the diversity of young neutron stars (\\citealt{alp01}; \\citealt*{ace13} and references therein). Unlike what was once thought---young radio pulsars like those in the Crab and Vela supernova remnants were prototypical of newborn neutron stars, it has been realized that there are classes of magnetars \\citep{wt06}, central compact objects (CCOs) in young supernova remnants \\citep*{pst04,del08}, and X-ray dim isolated neutron stars (XDINSs; \\citealt{tur09, mer11}). The latter three classes of young neutron stars are rather `quiet' at multiple wavelength regions from optical to radio: they generally do not have strong non-thermal emission and are not surrounded by any bright, pulsar-wind powered nebulae.  Neutron stars generally have non-thermal emission radiated from their magnetospheres, and sometimes a thermal component arising from their hot surfaces may be seen (e.g., \\citealt{bec09}; note that the thermal emission can be dominant in some cases). The non-thermal emission can be described by a power law with flux decreasing from X-ray to optical/IR wavelengths, while because of their $\\sim$10$^6$~K surface temperature, the thermal component's Rayleigh-Jeans tail may be detectable at ultraviolet/optical wavelengths (e.g., \\citealt{kap+11}). As a result, among $\\sim$2000 known neutron stars \\citep{man+05}, only over 20 neutron stars have been detected at optical/IR wavelengths. For a comparison, a disk would have thermal-like emission, and depending on temperature, which is often assumed to be a function of disk radius, the disk would be generally bright at optical/IR wavelengths \\citep{phn00}. Emission from putative disks would thus be distinguishable from that of neutron stars. In addition, if a debris disk consists of cold, $\\leq$100~K dust, it would also possibly be detectable at submillimeter (submm) and millimeter (mm) wavelengths \\citep{pc94}.  Given all these reasons, searches for disks around neutron stars have been carried out with different telescopes. The goal of the searches is to find thermal-like emission from neutron stars at wavelengths of from optical to submm/mm, which is not expected to be radiated from neutron stars themselves.  In this paper, we review the current status of searches and provide a summary of the main results. It should be note that since young neutron stars (for example the magnetars that are covered in this paper) may exhibit strong variability (e.g., \\citealt{kas07}), and hence (nearly) simultaneous. observations at multiple wavelengths are often required in order to identify the source of emission. ",
        "conclusions": "\\label{dis} Sufficient observational evidence has shown that debris disks are ubiquitous around from different main-sequence stars (e.g., see \\citealt{wya08} and references therein) to white dwarfs (e.g., \\citealt{fjz09}; \\citealt{xj12} and references therein). In main-sequence stars, debris disks may likely be the remnants of protoplanetary disks at early ages of $\\sim$10~Myr and then are replenished by planetesimal collisions. In white dwarfs, which like the neutron stars are also post main-sequence, compact stars, debris disks are thought to be formed from material produced by tidal disruption of planetary bodies \\citep{gra+90,jur03} that have survived through late phases of stellar evolution \\citep{ds02}. Comparing neutron stars to them, the major uncertainties would be the initial conditions and heating mechanisms if the current scenario of fallback disk formation is believed. For young neutron stars (magnetars, CCOs, and XDINSs) and regular radio pulsars, a debris disk would be the remnant of a fallback disk. The initial mass of the fallback disk and its interaction with a newly-born neutron star are quite uncertain, which would affect subsequent evolution and detectability. While dust grains around normal stars or white dwarfs are nicely heated to 100--1000~K temperature by ultraviolet/optical photons from host stars, revealing their existence, it is not clear how much a pulsar's particle wind can heat up a disk. Hopefully in the near future, with great capabilities of the next generation telescopes such as the thirty meter telescope and \\textit{James-Webb Space Telescope}, these uncertainties would be cleared out by detections or even sufficiently deep upper limits of a number of different classes of neutron stars.  In summary, deep searches for disks around different types of neutron stars have been carried out. For magnetars, 6 of them were found with NIR emission, which was seen (except for 4U~0142+61) to be related to magnetars' outburst activities. Further observations to identify the origin of the NIR emission is needed. Among the 6 magnetars with NIR counterparts, 4U~0142+61 has been found to have a MIR component in its optical and IR broad-band spectrum and it is likely indicative of a dust disk around this X-ray emitting neutron star. Deep MIR imaging of the magnetars 1E~2259+586 and 1E 1048.1$-$5937 detected the first source at MIR 4.5~$\\mu$m but did not detect the second one, and the results suggest a similar dust disk to that of 4U~0142+61 around the first source (although only on the basis of the MIR and NIR $K_s$ band detections) and exclude the existence of a similar disk around the second one. For CCOs and XDINSs, the deepest optical and NIR observations were carried out. The non-detections of them rule out accretion (from a disk) as the power source to produce the observed X-ray emission from the CCOs, while the upper limits on the XDINSs are not very constraining due to possibly large inner radii considered for their putative accretion disks. Deep MIR imaging of CCOs was also conducted, but the obtained upper limits can not conclusively determine whether or not dust disks exist. A few millisecond and middle-aged radio pulsars were searched. The discovery of a candidate MIR counterpart to the Vela pulsar, if confirmed, would provide a certain case for our understanding of disk heating and detectability of disks around regular pulsars. Extensive searches using WISE imaging data are underway, and hopefully the results would bring to light the general existence of debris disks and their evolution under pulsars' extreme environments.  We gratefully thank anonymous referees for constructive suggestions. This research was supported by the National Natural Science Foundation of China (11073042, 11373055) and the Strategic Priority Research Program ``The Emergence of Cosmological Structures\" of the Chinese Academy of Sciences (Grant No. XDB09000000).  Z.W. is a Research Fellow of the One-Hundred-Talents project of Chinese Academy of Sciences."
    },
    "1402/1402.0082_arXiv.txt": {
        "abstract": "The electromagnetic energy equation is analyzed term by term in a 3D simulation of kinetic reconnection previously reported by \\citet{vapirev2013formation}. The evolution presents the usual 2D-like topological structures caused by an initial perturbation independent of the third dimension. However, downstream of the reconnection site, where the jetting plasma encounters the yet unperturbed pre-existing plasma, a downstream front (DF) is formed and made unstable by the strong density gradient and the unfavorable local acceleration field. The energy exchange between plasma and fields is most intense at the instability, reaching several $\\unit{pW/m^3}$, alternating between load (energy going from fields to particles) and generator (energy going from particles to  fields) regions. Energy exchange is instead purely that of a load at the reconnection site itself in a region focused around the x-line and elongated along the separatrix surfaces. Poynting fluxes are generated at all energy exchange regions and travel away from the reconnection site  transporting an energy signal of the order of about $\\bfS \\approx 10^{-3} \\unit{W/m^2}$. ",
        "introduction": "Reconnection is one of the most studied processes capable of releasing magnetic energy into kinetic energy, in the form of flows and particle heating. The typical conditions for reconnection have a sheared magnetic field where across an interface at least one component of the magnetic field reverses sign. In that situation the field lines of opposite polarity break and reconnect in a new configuration. In the process the magnetic energy content is decreased in favor of energization of the particles, in the form of ordered flows or random particle heating.  The exact point where the magnetic field breaks needs not bear any direct link with the region where energy is in fact released. Traditionally, two competing scenarios have been guiding the discussion. In the Sweet-Parker mechanism \\citep{sweet,parker} a diffusion region is present around the point of topological line breakage and the energy dissipation in the form of Ohmic heating and plasma acceleration is assumed to take place in the diffusion layer. In competition, the Petschek \\citep{petschek} model sees standing slow shocks form an interface across which energy is released from the magnetic field to the particles.  More recently, the kinetic description has identified that aspects of both models are present in the real plasma (see \\citet{birn-priest} for a recent review). And real collisionless plasmas are kinetic of course. The electrons and ions become decoupled~\\cite{sonnerup1979solar,terasawa1983hall} and two layers reminiscent of the Sweet-Parker layer are present \\cite{laval1966instabilites,hesse1999diffusion,kuznetsova2000toward,gem1836}. An inner one has the electrons being accelerated and the outer one the ions. In the typical cartoon picture the two layers are nested boxes. Of course the cartoon is not real and the reality is more complex with the boxes extending along the separatrices, especially in presence of finite guide fields~\\cite{kleva,biskamp,ricciguide}.  But in kinetic reconnection, the energy release is not limited to the electron and ion diffusion regions proper. Other energy releases are possible in conduction with reconnection. A main process receiving substantial attention both from the theoretical simulation side and from direct in situ observation is that of the regions downstream of a reconnection site, where the plasma flow caused by reconnection slams into its surrounding plasma. In these regions the outflowing plasma and magnetic field act as a snowplow releasing its mass and energy against the surrounding pre-existing plasma and field.  Such fronts emanating from a reconnection site have been studied in recent  kinetic simulations \\citep{sitnov2009dipolarization} and confirmed by direct observational evidence has been obtained with data from the THEMIS mission~\\citep{runov2009themis}. Many effects relative to these fronts can be studied in 2D, but one key process requires a full 3D study:  the presence of an instability that  perturbs the downstream front (DF)~\\citep{nakamura2002interchange,guzdar2010simple}. Fluid studies have been used to capture the effect~\\cite{nakamura2002interchange,guzdar2010simple}. In the fluid case, the instability squires the nature of an interchange mode similar to the Rayleigh-Taylor instability. Across the downstream front the density increases substantially. The vertical magnetic field and the density pile up at the DF. The curvature of the field lines and the braking of the front by the momentum exchange with the yet unperturbed plasma leads to an effective acceleration pointing contrary to the DF speed. Such configuration is unstable to interchange modes: the higher density region is ahead of the front and the density gradient is therefore opposite to the direction of the acceleration~\\citep{nakamura2002interchange}.    A well know limitation of MHD in this circumstance is the inability to predict the fastest growing mode: the interchange instability has the same growth speed at any wavenumber $k$. This is circumvented in MHD by seeding the instability ad hoc~\\citep{guzdar2010simple} or self-contently as a consequence of other processes that produce the required seed~\\citep{lapenta2011self}.  Full kinetic studies are more suitable for modeling the process, but of course 3D kinetic simulations are much more computationally demanding, limiting the accessible domain size. In kinetic theory, extra physics is present to determine the scale of the process, predicting a fastest growing mode~\\citep{pritchett2010kinetic}. In particular the presence of a density (and pressure) gradient at the DF induces also drift waves and instabilities that modify the nature of the DF instability~\\citep{divin2013dipolarization}. Recently, \\citet{vapirev2013formation} reported a fully kinetic simulation of the development of the secondary instability in DF emerging from a reconnection region in 3D simulations of sufficient domain size to track the evolution for several ion skin depths.   Recent Cluster multispacecraft  observations provide direct evidence for the presence of an interchange instability at DFs\\citep{runov2012multipoint} at scales comparable with that observed in simulations~\\citep{guzdar2010simple,lapenta2011self,vapirev2013formation}.   In the present paper, we consider the issue of the energetic consequences of the DF instability. We track the electromagnetic energy through the system and observe its balance, as described by the electromagnetic energy equation, and its flow described by the Poynting vector.  The issue has been considered  from an observational angle in the recent work by \\citet{hamrin2012role}. The energy exchange between plasma and fields is measured by computing directly $\\bfJ \\cdot \\bfE$.  The measure is not easy, requiring the estimation of the current from the four Cluster spacecraft data \\citep{hamrin2011energy}. Nevertheless, it has been done and the published results are used here to compare with our simulation data.  Previous 2D studies have already investigated the energy exchanges. The energetics at the DF has been inferred from Cluster data by \\citet{huang2012kinetic}, finding an energy transfer from the fields to the plasma in DF, in agreement with  2D simulation results \\citep{sitnov2009dipolarization}. The 2D  simulations show a more intense electron contribution to the energy deposition near the x-point and a more intense ion deposition elsewhere~\\citep{sitnov2011onset}.  The extension to 3D can, and as shown below indeed does, explain several other features observed in data.  First, the observations \\citep{hamrin2012role} show that the energy exchange is exclusively from the fields to the plasma near the reconnection site, but as one considers Cluster crossings closer to the Earth the presence of energy exchanges in both directions is found. The condition where energy is going to the plasma from the fields is called a load, with circuit terminology. The opposite situation of energy being transferred by the particles to the field is called a generator. The presence of transfer of energy in both directions is shown also in MHD models~\\citep{birn2005energy}, with loads concentrated in the mid-nigh region and generator regions located in the flanks where the cross-tail current is diverted to the field-aligned currents of the substorm current wedge. The mechanism at play here is  different and due to the development of the DF instability.  Second, there is observational evidence for an important role of waves  generated at the DF \\citep{marghitu2006experimental} where  generator regions  radiate electromagnetic energy. The suggestion is made that the energy goes into kinetic \\alf waves.   The study for the  energy fluxes from Cluster observations shows that the dominant component of the energy flux is ion enthalpy flux, with smaller contributions from the electron enthalpy and heat flux and the ion kinetic energy flux\\citep{angelopoulos2002plasma,eastwood2013energy}. The Poynting flux is a minority contribution but it is not negligible, and in certain parts of the ion diffusion region the Poynting flux in fact dominates\\cite{eastwood2013energy}  The present description revisits these issues using 3D fully kinetic simulations, measuring directly the energy balance and reproducing many of the observational features outlined above. In particular, we observe that electromagnetic energy is converted to plasma energy at the reconnection site in a region very elongated along the separatrices. At the DF, we find the instability to convert about one order of magnitude more energy than at the reconnection site itself. But this energy is alternating between generator and load regions, as in the observations mentioned above.  Significant regions of intense Poynting flux emerges from both regions of energy exchange. At the reconnection site, the energy flux is again of definite sign, but at the DF alternates in sign as the source causing it.  The remainder of the paper is organized as follows: Section 2 reports the details of the simulation approach followed and gives an overall view of he processes developing in the simulation. Section 3 considers the energy balance equation for the electromagnetic energy and investigates each term in turn. Section 4 provides a summary and an interpretation of the results of the energy analysis making a direct link with observations. The final overview is provided in Fig.~\\ref{scenario} that summarizes pictorially the findings of the present investigation. ",
        "conclusions": ""
    },
    "1402/1402.2614_arXiv.txt": {
        "abstract": "The classical picture of a star-forming filament is a near-equilibrium structure, with collapse dependent on its gravitational criticality. Recent observations have complicated this picture, revealing filaments as a mess of apparently interacting subfilaments, with transsonic internal velocity dispersions and mildly supersonic intra-subfilament dispersions. How structures like this form is unresolved. Here we study the velocity structure of filamentary regions in a simulation of a turbulent molecular cloud. We present two main findings: first, the observed complex velocity features in filaments arise naturally in self gravitating hydrodynamic simulations of turbulent clouds without the need for magnetic or other effects. Second, a region that is filamentary only in projection and is in fact made of spatially distinct features can displays these same velocity characteristics. The fact that these disjoint structures can masquerade as coherent filaments in both projection and velocity diagnostics highlights the need to continue developing sophisticated filamentary analysis techniques for star formation observations. ",
        "introduction": "Observations with the Herschel satellite have revealed that the backbone of molecular clouds is a compex network of connecting and interacting filaments \\citep[e.g.][]{2010A&A...518L.102A,2010A&A...518L.100M,2011A&A...529L...6A,2012A&A...540L..11S}. The dynamics of this web of dense molecular gas not only determines the evolution and stability of molecular clouds, but also regulates their condensation into stars. Molecular cloud cores and single low-mass stars are almost always found in filaments, often aligned like pearls on a string \\citep{2002ApJ...578..914H,2008ApJ...672..410L,2010A&A...518L.102A}. This can be interpreted as a result of the gravitational instability of a supercritical filamentary section \\citep{1997ApJ...480..681I,2011A&A...533A..34H}. Where filaments intersect, more massive hubs form \\citep{2009ApJ...700.1609M,2013ApJ...764..140M} that later on could become the progenitors of star clusters.   \\citet{2010ApJ...724..687L} \\citep[see also e.g.][]{2007ApJ...666..982E} demonstrated that the fraction of gas in the dense molecular web is of order 10 percent of the total molecular mass. Interestingly, the mass fraction of protostars to dense ($n > 10^4$ cm$^{-3}$) molecular gas is also a constant of the same order. \\citet{2013ApJ...773...48B} showed that this requires new filamentary segments to continuously form from the diffuse intra-filament medium on a gravitational collapse timescale.  Classically, gas filaments have been treated as cylinders of gas in hydrostatic equilibrium \\citep{1964ApJ...140.1529O,2012A&A...542A..77F,2013A&A...558A..27R}. It is not clear, however, how such quiescent structures could form in the turbulent environment of a molecular cloud \\citep[e.g.][]{2010A&A...520A..17K,2013ApJ...769..115H}. In addition, millimeter line studies indicate that filaments have intrinsic, super-thermal linewidths \\citep{2013A&A...553A.119A}.  More recently, observations by \\citet{2013A&A...554A..55H} revealed that filaments are often compact bundles of thin spaghetti-like subfilaments. They presented observations of a prominent filamentary feature in Taurus, dominated by the L1495 cloud \\citep{1962ApJS....7....1L} and several dark patches \\citep{1927cdos.book.....B}. They observed the region in the moderate density tracer C$^{18}$O, obtaining spectra along a $\\sim 10$ pc length of the region. Analysing the richest $\\sim 3$ pc section of the filament in the resultant position--position--velocity space, they found the intriguing result that the gas along the ridge is organized in velocity-coherent filamentary structures, with typical lengths $\\sim0.5$ pc. Each filament is internally subsonic or transsonic, though the collection of filaments is characterised by a mildly supersonic interfilamentary dispersion of $\\sim 0.5$ km s$^{-1}$. They describe the collection of velocity detections as \"elongated groups that lie at different `heights' (velocities) and present smooth and often oscillatory patterns\". Multiple velocity components along a single line of sight have also been observed in Serpens South \\citep{2013ApJ...778...34T}; this feature may therefore be common for many young star forming sites.  The complexity of Hacar et al.'s position--position--velocity data, exemplified by their figure 9, invites theoretical and numerical exploration. In particular the apparent organization of filaments into bundles tends to bring to mind magnetic fields or other relatively complex physics. In this paper we explore whether this velocity signal is present in simulated molecular clouds with a more minimal set of physics, including only gravity and hydrodynamic forces acting on the initial turbulence.  \\begin{figure*} \\begin{center} \\includegraphics[width=2\\columnwidth]{figure1.pdf} \\caption{\\label{finderimage} Surface density projections through the simulated volume at two times. The dashed cyan lines mark the boundary of the periodic simulation domain. We calculated the surfaced density using only the gas between $10^3 < n/{\\rm cm}^{-3} < 10^{4.5}$, and converted this to an approximation of an optically thin C$^{18}$O observation. The magenta boxes delineate the filaments we focus on in this paper. Sink particles are shown as orange dots.} \\end{center} \\end{figure*} ",
        "conclusions": "We simulated a 10pc region of a turbulent molecular cloud, allowing it to evolve under self gravity for 1.25 Myr. While the global structure of the region is still dominated by its turbulent structure, the denser gas has had time to become gravitationally organized. We treated the simulation as an observer might observe the sky; converting the density to an approximate line intensity, selecting filamentary regions for further study, and analysing the line-of-sight velocity information in these regions.  Our first main finding is that velocity characteristics very similar to those observed by \\citet{2013A&A...554A..55H} form naturally in such a turbulent setup. Individually bound subfilaments display approximately sonic or subsonic dispersions, while the agglomerations that make up the larger filaments have transsonic to mildly supersonic relative motions. While we will further study the detailed evolution of these structures in future work, we speculate that this velocity structure is a relic of the supersonic turbulence that is generally taken as the initial conditions of star formation. The substructured (both in space and velocity) filaments appear without the need for magnetic fields, which are the physical mechanism that immediately spring to mind when considering a filament composed of a bundle of subfilaments.  Secondly, the same velocity structures that appear in spatially coherent filaments can show up when observing a faux-filament. Based on the line-of-sight velocity information, there is not a clear case to be made that our filament B (particularly the section with $L > 2$pc; see Figure \\ref{filvels1}) is different from filaments C1 or D1. Examining the third dimension reveals it to in fact be composed of widely separated dense regions (Figure \\ref{filaments3D}). Apparent velocity coherence similar to that seen in our most linear structures is evidently not enough to diagnose a true filament. The presence of this imposter filament in our data, not merely in projected density but also in the line-of-sight velocity, highlights the need for the continued development of sophisticated filament diagnostics, using both simulation and observation and both spatial and velocity data, in order to interpret not only existing observations of star forming filaments but also the imminent onslaught of data from ALMA."
    },
    "1402/1402.6684.txt": {
        "abstract": "The Automated Planet Finder (APF) is a facility purpose-built for the discovery and characterization of extrasolar planets through high cadence Doppler velocimetry of the reflex barycentric accelerations of their host stars. Located atop Mt. Hamilton, the APF facility consists of a 2.4-m telescope and its Levy Spectrometer, an optical echelle spectrometer optimized for precision Doppler velocimetry. APF features a fixed format spectral range from 374 nm - 970 nm, and delivers a ``Throughput\" (resolution $*$ slit width product) of 114,000 arc-seconds, with spectral resolutions up to 150,000. Overall system efficiency (fraction of photons incident on the primary mirror that are detected by the science CCD) on blaze at 560 nm in planet-hunting mode is 15\\%. First-light tests on the RV standard stars HD 185144 and HD 9407 demonstrate sub-\\ms\\ precision (RMS per observation) held over a 3-month period. This paper reviews the basic features of the telescope, dome, and spectrometer, and gives a brief summary of first-light performance. ",
        "introduction": "Around the turn of the millennium, the California-Carnegie Exoplanet Team launched a project to construct a dedicated ground-based precision radial velocity facility with which to find planets in the liquid water habitable zone \\citep{kas93, kas13} around low mass stars. An early discussion of the prevalence and possibility of potentially habitable planets around M dwarfs can be found in the discovery paper for GJ 876b, the first known planet around an M dwarf \\citep{mar98}. Such planets have periods in the 20-60 day regime that are quite hard to capture from a conventionally shared telescope that is scheduled around lunation phases. One needs telescope access throughout the lunar month to overcome aliasing and phase coverage problems for periods near the lunar month or integral multiples thereof. Also, one needs much higher cadence than is typically obtainable with a shared large telescope such as Keck in order to acquire enough data points to realize substantial $\\sqrt N$ gains in signal-to-noise to tease out  the extremely weak signals of Earth-sized planets in habitable zones from a limiting background of essentially random stellar noise. Successful observational and modeling strategies for optimizing exoplanet detection against a limiting background of stellar jitter noise have been presented by \\cite{dum11a, dum11b, tuo12} and others. And, as this facility was to be operated every night of the year, we also wished to automate the telescope, and hence named it the Automated Planet Finder or APF.  The APF project got its official start with cornerstone funding obtained in the form of a \\$6.4 million ear-mark in the 2002 Defense Appropriations Bill through the U.S. Naval Observatory (USNO). For that amount, the original intent was to simply purchase an existing clone of the four 1.8-m telescopes from Electro-Optical Systems Technologies (EOST) of Tucson, AZ that were originally slated to be used as outrigger telescopes for the Keck twin-telescope interferometer. That clone was in storage in EOST's Tucson warehouse. The plan was to house this 1.8-m telescope in a commercial dome sited in the parking lot adjacent to the Shane 3-m telescope's dome and fiber-feed its output into the existing Hamilton spectrometer \\citep{vog87} of the Shane 3-m. However, an additional \\$1.81M NASA grant was obtained that allowed the project to be super-sized up to a 2.4-m EOST telescope and dedicated new spectrometer, in an IceStorm-2 dome manufactured by EOST's parent company Electro-Optical Systems (EOS) of Queanbeyan, Australia. The new dome was sited next to the Lick 20\" Astrograph telescope atop Mt. Hamilton. The total cost-to-completion of the APF project was \\$12.37 million.  Work on the project began in 2004 under supervision of the UCO/Lick Labs in Santa Cruz with S. Vogt as Principal Investigator, G. Marcy and D. Fischer as Co-I's, and M. Radovan as Project Manager and Principal Engineer. A large team of technical support personnel from UCO/Lick  provided the requisite site engineering and construction management, as well as providing major engineering support to EOS for the dome construction and commissioning, to EOST for telescope construction and commissioning, and for building the Levy spectrometer. In 2007 the California-Carnegie Exoplanet team split into two teams, the California Planet Survey (CPS) team led by G. Marcy at UC. Berkeley, and the Lick-Carnegie Exoplanet Survey (LCES) team co-led by R. Butler at DTM and S. Vogt at UCO/Lick. Although the APF project was initiated under the California-Carnegie Exoplanet team, oversight and management of the telescope and dome contracts, fabrication of the spectrometer, and commissioning of the overall facility were carried out solely by the LCES team and technical staff of UCO/Lick Observatory.  The project was initially expected to require only 2-3 years to complete and the contracts called for the delivery of a turn-key telescope and integrated dome facility. Getting to the June 2009 acceptance sign-off involved an arduous six-year slog, with major support on virtually all fronts by UCO/Lick technical staff to the prime contractors for both the telescope and dome. Both the telescope and dome systems had serious issues that required extensive assistance from Lick technical staff to resolve. Among these were UCO/Lick having to make a new M2 to correct a serious back-focal-distance error, providing major assistance in the polishing and figuring of M1, recoating of M3, and providing also the coating for M1. Since then June 2009 acceptance sign-off, UCO/Lick staff have struggled to overcome numerous post-acceptance issues with both the dome and telescope, a task made substantially more difficult by the fact that, while some of the critical documentation (software and firmware) is locked away in an escrow account until January 29, 2018, other critical firmware and board schematics were never delivered and were not included in the escrow deposit. ",
        "conclusions": "This paper presents a brief overview of the newly-commissioned Automated Planet Finder or APF. At the time of this writing, the facility is still in the final commissioning and testing phases. The facility is currently capable of working through an entire night autonomously, and has been mostly operating robotically since Jan. 1, 2014.  Unlike other highly successful precision RV facilities such as HARPS, HARPS-N, and SOPHIE-HR, APF does not attempt to achieve precision through absolute stability of the instrument and the use of an image scrambler. Rather, an Iodine cell is used to provide the precision velocity reference, thereby greatly easing demands on long-term instrumental stability. Nevertheless, the Levy spectrometer on APF incorporates extensive passive athermalization in its optomechanical design to help enhance stability of the PSF and spectral position on the CCD. This passive athermalization, together with a temperature-controlled insulated housing, keeps the spectrum position stable on the CCD to typically under 0.2 pixels over at least several-months time scales, at or a bit less than the position drifts due to barometric pressure changes. Unlike HARPS, where the entire spectrometer is enclosed in a vacuum chamber, or CHIRON \\citep{tok13} and PFS where the echelle is enclosed in a vacuum chamber, no attempt was made with the Levy to control or stabilize atmospheric pressure. Such sub-pixel spectral shifts due to atmospheric pressure changes have no discernible deleterious affect on the resultant precision, and are faithfully removed by the Iodine reference lines. Accurate guiding is also crucial to obtaining high RV precision. Guiding on APF is done by picking off 4\\% of the incoming starlight with a beamsplitter and presenting a symmetric unvignetted seeing disk to the CCDTV for more precise guiding. Tight closed loop tracking specs on the servo-driven telescope also help to further stabilize guiding.  Since the APF has forgone the use of fiber scrambling and/or image slicing (both of which can produce substantial losses if not extremely well executed technically), it achieves rather high overall optical efficiency. Under typical exoplanet hunting conditions, the APF facility delivers a peak efficiency (fraction of photons hitting the primary mirror that are detected by the science CCD) of about 15\\%, with a typical spectral resolution of 110,000. At the same time, despite the lack of an image scrambler to stabilize the PSF, observations of the RV-null stars Sigma Dra (HD 185144) and HD 9407 over a period of three months in summer of 2013 demonstrated sub-\\ms\\ (RMS per point) instrumental precision.  The APF's Levy produces a fixed spectral coverage from 374-900 nm, much wider than the 490-600 nm Iodine region. As such, it should be useful for a wide range of conventional stellar spectroscopy programs other than precision RV's. While its nominal dispersion and resolution are quite high by traditional standards for fainter object work, it may prove useful for target of opportunity (ToO) observations in support of bright supernovae and GRB follow-up. APF's  guider currently can reach V=15, with extension to at least V=18.5 when using the charge multiplication mode on the auto guider CCDTV. But for it's primary mission, high precision RV work, APF's high efficiency coupled with sub-\\ms\\ precision and dedicated nightly cadence abilities, should make it a valuable contributor over the coming years to exoplanet discovery and characterization."
    },
    "1402/1402.4803_arXiv.txt": {
        "abstract": "Mergers of binary neutron stars (NSs) usually result in the formation of a hypermassive neutron star (HMNS).  Whether- and when this remnant collapses to a black hole (BH) depends primarily on the equation of state and on angular momentum transport processes, both of which are uncertain.  Here we show that the lifetime of the merger remnant may be directly imprinted in the radioactively powered \\emph{kilonova} emission following the merger.  We employ axisymmetric, time-dependent hydrodynamic simulations of remnant accretion disks orbiting a HMNS of variable lifetime, and characterize the effect of this delay to BH formation on the disk wind ejecta. When BH formation is relatively prompt ($\\lesssim 100$ ms), outflows from the disk are sufficiently neutron rich to form heavy $r$-process elements, resulting in $\\sim$week-long emission with a spectral peak in the near-infrared (NIR), similar to that produced by the dynamical ejecta.  In contrast, delayed BH formation allows neutrinos from the HMNS to raise the electron fraction in the polar direction to values such that potentially \\emph{Lanthanide-free} outflows are generated.  The lower opacity would produce a brighter, bluer, and shorter-lived $\\sim$ day-long emission (a `blue bump') prior to the late NIR peak from the dynamical ejecta and equatorial wind. This new diagnostic of BH formation should be useful for events with a signal to noise lower than that required for direct detection of gravitational waveform signatures. ",
        "introduction": "Mergers of binary neutron stars (NSs) (hereafter `neutron star mergers', or NSMs) are the primary source of gravitational waves (GW) for upcoming ground-based interferometric detectors such as Advanced LIGO and Virgo (\\citealt{Abadie+10}).  They are also promising central engines for short-duration gamma-ray bursts (GRBs; \\citealt{Paczynski86,Eichler+89}; see \\citealt{Berger13} for a recent review).  General relativistic simulations of NSMs show that the merger process can result in two qualitatively different outcomes, depending primarily on the total mass of the binary $M_{\\rm t}$.  If $M_{\\rm t}$ exceeds a critical value $M_{\\rm c}$, then the massive object produced by the merger collapses to a black hole (BH) on the dynamical time ($\\sim$few ~ms, e.g., \\citealt{sekiguchi2011}). On the other hand, if $M_{\\rm t} < M_{\\rm c}$ then the merger product is at least temporarily supported against gravitational collapse by differential rotation and/or thermal pressure. This meta-stable compact object is usually called a \\emph{hypermassive} NS (HMNS; e.g. \\citealt{kaplan2013}).  The value of $M_{\\rm c}$ depends on the uncertain equation of state of (EoS) of nuclear matter. The recent discovery of massive $\\sim 2M_{\\odot}$ NSs (\\citealt{Demorest+10,Antoniadis+13}) excludes a soft EoS, placing a lower limit of $M_{\\rm c} \\gtrsim 2.6-2.8M_{\\odot}$ (\\citealt{Hotokezaka+13}; \\citealt{Bauswein+13}).  It thus appears likely that the `canonical' 1.4 + 1.4 $M_{\\odot}$ binary merger goes through a HMNS phase.   \\begin{figure*} \\includegraphics*[width=1.5\\columnwidth]{f1.eps} \\caption{ Relation between the observed kilonova and the properties of the ejecta that powers it. Material ejected dynamically in the equatorial plane is highly neutron rich ($Y_e < 0.1$), producing heavy r-process elements that include Lanthanides. This results in emission that peaks in the near-infrared and lasts for $\\sim 1$~week (`late red bump') due to the high opacity. Outflows from the remnant disk are more isotropic and also contribute to the kilonova. If the HMNS is long-lived, then neutrino irradiation can increase $Y_e$ to a high enough value ($Y_e\\sim 0.4$) that no Lanthanides are formed, resulting in emission peaking at optical wavelengths (`early blue bump'). If BH formation is prompt, outflows from the disk remain neutron rich, and their contribution is qualitatively similar to that of the dynamical ejecta. } \\label{f:cartoon} \\end{figure*}  When a HMNS does form, its lifetime before collapsing into a BH depends on the timescale for thermal energy loss via neutrino emission (e.g., \\citealt{ruffert1999,Paschalidis+12,Galeazzi+13}) and the efficacy of angular momentum transport via gravitational waves and magnetohydrodynamic stresses (e.g., \\citealt{Duez+06,Stephens+08,Siegel+13}). For particularly nearby mergers, oscillations excited in the HMNS may be detectable in the GW strain data (\\citealt{Shibata05}; \\citealt{Bauswein+12}; \\citealt{Hotokezaka+13c}).  However, the subsequent ring down phase following BH formation is unlikely to be detected by the initial generation of advanced detectors.  Fortunately, NSMs are also accompanied by coincident electromagnetic (EM) signals that inform physical processes at work during the merger (e.g.~\\citealt{Metzger&Berger12}; \\citealt{Kelley+13}; \\citealt{Piran+13}). One such counterpart is a thermal IR/optical transient powered by the radioactive decay of heavy elements synthesized in the merger ejecta (a `kilonova'; \\citealt{Li&Paczynski98}; \\citealt{Metzger+10}; \\citealt{Roberts+11}; \\citealt{Goriely+11}; \\citealt{Piran+13}; \\citealt{Grossman+13}; \\citealt{Tanaka+14}).  Kilonovae are particularly promising EM counterparts because (1) their generation is relatively robust, requiring only a modest amount of unbound ejecta; (2) their signal is independent of the existence of a dense surrounding external medium; and (3) unlike a GRB, kilonovae are relatively isotropic. A candidate kilonova was recently detected following the GRB 130603B (\\citealt{Tanvir+13}; \\citealt{Berger+13}; ).  If the merger ejecta is sufficiently neutron-rich for $r$-process nucleosynthesis to reach the Lanthanides ($A \\gtrsim 139$), the optical opacity becomes much higher than that of iron-group elements \\citep{Kasen+13}, resulting in emission that is redder, dimmer, and more slowly evolving (\\citealt{Barnes&Kasen13,tanaka2013}).  Although such unusually red colors may be beneficial in distinguishing NSM transients from unrelated astrophysical sources, the current lack of sensitive wide field infrared telescopes could make EM follow-up across the large sky error regions provided by Advanced LIGO/Virgo even more challenging (e.g.~\\citealt{Nissanke+13}; \\citealt{Metzger+13}; \\citealt{Hanna+13}; \\citealt{Kasliwal&Nissanke13}).  The matter ejected dynamically following a NSM is likely to be sufficiently neutron-rich (as quantified by the electron fraction $Y_e~\\lesssim 0.3$) to produce a red kilonova (e.g., \\citealt{Rosswog05,Duez+09,bauswein2013}). Dynamical expulsion is not the only source of ejecta, however.  A robust consequence of the merger process is the formation of a remnant torus surrounding the central HMNS.  Outflows from this accretion disk over longer, viscous timescales also contribute to the merger ejecta (e.g.,~\\citealt{Surman+08,Metzger+08,Metzger+09a,Lee+09,Dessart+09,Wanajo&Janka12}). The more isotropic geometry of disk winds suggests that they may contribute a distinct component to the kilonova light curve for most viewing angles \\citep{Barnes&Kasen13, Grossman+13}.  \\citet[ hereafter FM13]{FM13} calculated the viscous evolution of remnant BH accretion disks formed in NSMs using two-dimensional, time-dependent hydrodynamical simulations.  Over several viscous times, FM13 found that a fraction $\\sim$several percent of the initial disk mass is ejected as a moderately neutron-rich wind ($Y_e \\sim 0.2$) powered by viscous heating and nuclear recombination.  Although the higher entropy of the outflow as compared to the dynamical ejecta results in subtle differences in composition (e.g. a small quantity of helium), the disk outflows likely produce Lathanide elements with sufficient abundance to result in a similarly red kilonova as with the dynamical ejecta.   FM13 included the effects of self-irradiation by neutrinos on the dynamics and composition of the disk. Due to the relatively low accretion rate and radiative efficiency at the time of the peak outflow, neutrino absorption had a sub-dominant contribution to the disk evolution. This hierarchy is important because a large neutrino flux tends to drive $Y_e$ to a value higher than that in the disk midplane (e.g.~\\citealt{Surman+08}; \\citealt{Metzger+08}; \\citealt{Surman+13}).  If neutrino irradiation is sufficient to drive $Y_e \\gtrsim 0.3-0.4$, the nuclear composition of the disk outflows would be significantly altered, resulting in a distinct additional component visible in the kilonova emission.  By ignoring the influence of a central HMNS, FM13 implicitly assumed a scenario in which BH formation was prompt or the HMNS lifetime very short. Here we extend the study of FM13 to include the effects of neutrino irradiation from a long-lived HMNS. As we will show, the much larger neutrino luminosity of the HMNS has a profound effect on the quantity and composition of the disk outflows, allowing a direct imprint of the HMNS lifetime on the kilonova (Figure~\\ref{f:cartoon}). As in FM13, our study includes many approximations that enable us to follow the secular evolution of the system. We focus here on exploring the main differences introduced by the presence of a HMNS, and leave more extensive parameter space studies or realistic computations for future work.  The paper is organized as follows.  In $\\S\\ref{s:model}$ we describe the numerical model employed. Our results are presented in $\\S\\ref{s:results}$, separated into dynamics of the outflow (\\S\\ref{s:evolution}) and composition (\\S\\ref{s:composition}). A summary and discussion follows in $\\S\\ref{s:discussion}$. Appendix~\\ref{s:neutrino_details} describes in more detail the upgrades to the neutrino physics implementation relative to that of FM13. ",
        "conclusions": "\\label{s:conclusions} \\label{s:discussion}  We have explored the effects of a hypermassive neutron star (HMNS) on the long-term evolution of remnant accretion disks formed in neutron star binary mergers. Our main results can be summarized as follows: \\newline  \\noindent 1. -- A long-lived HMNS results in the ejection of a significant fraction of the disk over a timescale of $\\sim 1$~s. The amount of mass increases monotonically with HMNS lifetime. This enhanced mass loss, up to a factor $\\gtrsim 10$ relative to prompt BH formation, results from enhanced neutrino heating and a reflecting inner boundary. \\newline  \\noindent 2. -- The composition of the ejecta is latitude-dependent. Material within an angle $\\sim 30^\\circ$ of the polar axis is strongly irradiated by the HMNS. If the neutrino and antineutrino luminosities are similar (as we have assumed), the electron fraction of the outflow can be raised to values where Lanthanides are not produced. \\newline  \\noindent 3. -- Material ejected equatorially is still expected to produce a strong r-process, similar to that of the dynamical ejecta, although the detailed composition is dependent on the HMNS lifetime. \\newline  The criterion on the neutron star lifetime to appreciably change the electron fraction of the polar outflow ($t_{\\rm ns} \\gtrsim 100$ ms) is obviously satisfied if the remnant mass is below the maximum mass of a cold neutron star, or if the remnant is stabilized by solid body rotation (a {\\it supramassive} NS), the latter of which is only removed on much longer timescales via e.g. magnetic dipole spin-down. The $t_{\\rm ns} \\gtrsim 100$ ms condition is also likely satisfied if the HMNS is supported by thermal pressure (e.g.~\\citealt{Bauswein+10}; \\citealt{Paschalidis+12}; \\citealt{kaplan2013}), as the latter is removed on the HMNS cooling timescale, which is typically on the order of $\\sim$ seconds (e.g.~\\citealt{pons1999}).  On the other hand, support via differential rotation may not last 100 ms, as the latter can be efficiently removed by magnetic stresses: the growth rate of the MRI is as short as milliseconds.  Whether a blue bump indeed develops out of the polar outflow in the case of a long-lived HMNS will depend on a number of factors. First, enough mass needs to be ejected so that the contribution to the lightcurve from radioactive decay becomes detectable.  Second, the level of irradiation and the ratio between neutrino and antineutrino luminosities must be sufficient to raise $Y_e$ to values close to $0.5$.  The results of \\citet{Dessart+09}, who employ a much more realistic treatment of neutrino transport but did not include the viscous evolution, seem to align with our findings. The ejecta mass required for a detectable signal depends on how much radioactive heating is supplied by synthesized elements, which are lighter than the $r$-process nuclei but potentially more neutron-rich than $^{56}$Ni (e.g.~\\citealt{Grossman+13}).  A final requirement for producing a blue bump is that the high $Y_e$ (Lanthanide-free) material must be ejected first, such that it resides {\\it exterior} to any lower $Y_e$ (Lanthanide-rich) material that would otherwise block its emission.  Our calculations support this requirement as well: after rising quickly, $Y_e$ of the polar ejecta decreases monotonically with time (Fig.~\\ref{f:ye_average}). This is a direct consequence of the decrease in the neutrino luminosities with time, either smoothly as in the case of a long-lived HMNS (Appendix~\\ref{s:neutrino_details}), or via a sudden drop when a BH forms. Future work will explore in more detail the observational consequences of this bimodal outflow.  Figure~\\ref{f:summary_plot} provides estimates of various quantities relevant to the kilonova emission derived from our calculations, as a function of the HMNS collapse time.  The peak luminosities and time to peak are estimated using the `Arnett rule' for a kilonova (e.g., \\citealt{Li&Paczynski98,Metzger+10}), \\begin{eqnarray} \\label{eq:L_peak} L_{\\rm peak} & \\simeq & 4.3\\times 10^{41}\\textrm{ erg s}^{-1}\\left(\\frac{f}{3\\times 10^{-6}} \\right) \\left(\\frac{v_r}{0.1c} \\right)^{1/2}\\nonumber\\\\ &&\\qquad\\qquad\\qquad\\qquad \\times \\left(\\frac{M_{\\rm ej}}{0.01M_\\sun} \\right)^{1/2}\\kappa^{-1/2}\\\\ \\label{eq:t_peak} t_{\\rm peak} & \\simeq & 1.4\\textrm{ d } \\left(\\frac{v_r}{0.1c} \\right)^{-1/2} \\left(\\frac{M_{\\rm ej}}{0.01M_\\sun} \\right)^{1/2}\\kappa^{1/2}, \\end{eqnarray} where $f$ is a factor quantifying radioactive energy deposition, and $\\kappa$ is the opacity of the material in units of cm$^2$~g$^{-1}$. In making Figure~\\ref{f:summary_plot}, we have used $\\kappa = 10$~cm$^2$~g$^{-1}$ for the equatorial material (Lanthanide-dominated), and $\\kappa = 1$~cm$^2$~g$^{-1}$ for the polar material (Lanthanide-free), see \\citet{Kasen+13}. The velocity is the mass-flux weighted value at $r = 10^9$~cm considering only unbound material.  \\begin{figure} \\begin{overpic}[width=\\columnwidth]{f6.eps} \\put(2.5,93.3){\\tiny $\\odot$} \\end{overpic} \\caption{Kilonova properties as a function of HMNS lifetime. Shown are the ejecta mass in unbound material (a), the mass-flux weighted velocity at $10^9$~cm (b), the mass-flux weighted electron fraction (c), the peak luminosity (d) (eq.~[\\ref{eq:L_peak}]), and the time to peak (e) (eq.~[\\ref{eq:t_peak}]). Red squares and blue dots indicate equatorial and polar material. The opacity is assumed to be Lanthanide-dominated for the former and Lanthanide-free for the latter. Note that these numbers underestimate the `red' component since we have not included contributions from the dynamical ejecta. } \\label{f:summary_plot} \\end{figure}  Most of the resulting kilonova properties are a monotonic function of the HMNS lifetime. Peak luminosities increase from $\\sim 10^{40}$ to $\\sim 10^{41}$~erg~s$^{-1}$ for the red component, while the blue component, when present, is brighter by a factor $\\sim 2$. Similarly, the peak times for the red component range from about a week to a month, increasing with longer HMNS lifetime due to the larger ejected mass. The blue component can last from a few- to several days. While the blue component is faster with longer HMNS lifetime, the average velocity of the red component saturates at $\\sim 0.05c$.  The large amount of ejecta mass found by our calculations suggests that outflows from the disk could easily overwhelm that from the dynamical ejecta. One implication of this result relates to the Galactic production of $r$-process elements.  FM13 estimated that ejection of $\\sim 10\\%$ of the disk mass would contribute with $\\sim 20\\%$ of the production rate of elements with $A\\gtrsim 130$ assuming reasonable values for the disk mass and neutron star merger rate. The increase in the ejected fraction of the disk by a factor of several  relative to the prompt BH case implies that disks with long-lived HMNS could become a dominant contribution to the galactic $r$-process element production (\\citealt{Freiburghaus+99}; \\citealt{Rosswog+13}; \\citealt{Piran+14}). This is the case even when the polar outflow is Lanthanide free (Fig~\\ref{f:histogram}).  The large ejecta masses we infer for a moderately long-lived HMNS ($t_{\\rm ns} = 100$ ms) may also help alleviate the tension between large ejecta mass $\\sim 3\\times 10^{-2}M_{\\odot}$ required to fit the NIR excess observed following GRB 130603B with current models (\\citealt{tanaka2013}; \\citealt{Piran+14}), without the need to invoke less likely scenarios such as the merger of a NS with a low mass BH.  If the long-lived HMNS is magnetized, its rotational spin-down could also power the excess X-rays observed following this event (\\citealt{Fong+14}; \\citealt{Metzger&Piro13}; \\citealt{Fan+13}).  Given that early blue emission appears to require the presence of a long-lived HMNS, detection of such a component in future events provides a relatively clean way to rule out a NS-BH merger (although the {\\it absence} of early blue emission would not rule out a NS-NS merger).  Our models include many approximations in order to make the evolution to a time $\\sim 10$ s computationally feasible.  In addition, we have focused here on the key differences introduced by the HMNS applied to models with a particular choice of parameters. Much more work remains in order to make reliable predictions for kilonovae emission. In addition to a more extensive exploration of parameter space,  models with realistic angular momentum -  and neutrino transport will be needed."
    },
    "1402/1402.5752_arXiv.txt": {
        "abstract": "We analyze the recent released $HST$/WFC3 IR images in the GOODS-N region to study the formation and evolution of Quiescent galaxies (QGs). After examining the reliability with artificial galaxies, we obtain the morphological parameters with S\\'ersic profile of 299 QGs and 1,083 star-forming galaxies (SFGs) at $z\\sim0.5$--3.0, finding the evolution of $r_\\mathrm{e}$ and $n$ of massive ($M_*\\geq10^{10.5}$~M$_\\sun$) QGs while weaker evolution of SFGs and less massive ($M_*<10^{10.5}$~M$_\\sun$) QGs. The regression of the size evolution of massive QGs follows $r_\\mathrm{e}\\propto (1+z)^{-\\alpha_{r_\\mathrm{e}}}$ with $\\alpha_{r_\\mathrm{e}}=1.06\\pm0.19$ (a factor of $\\sim2.2$ increase from $z\\sim2.5$ to $\\sim0.5$), which is consistent with the general picture of the significant size growth. For the further understanding of the evolution scenario, we study the evolution of S\\'ersic index, $n$, and find that of massive QGs to significantly evolve as $n\\propto(1+z)^{-\\alpha_n}$ with $\\alpha_n=0.74\\pm0.23$ ($n\\sim1$ at $z\\sim2.5$ to $n\\sim4$ at $z\\sim0.5$), while those of the other populations are unchanged ($n\\sim1$) over the redshift range. The results in the present study are consistent with both of observation and numerical simulations, where gas-poor minor merger is believed to be the main evolution scenario. By taking account of the connection with less massive QGs and SFGs, we discuss the formation and evolution of the massive QGs over{\\it ``Cosmic High Noon\"}, or the peak of star-formation in the universe. ",
        "introduction": "The study of the high-redshift (high-$z$) early-type galaxies (ETGs) provides us clues to understanding the formation and evolution of massive galaxies in the local universe. Their star-formation activity peaked during the cosmological epoch at $1 < z < 3$ (e.g., Dickinson et al.~\\citeyear{dickinson03}; Heavens et al.~\\citeyear{heavens04}; Papovich et al.~\\citeyear{papovich06}; Hopkins \\& Beacom~\\citeyear{hopkins06}) and galaxy morphologies have changed dramatically (Kajisawa \\& Yamada~\\citeyear{kajisawa06}). For galaxy sizes, many studies have corroborated that massive galaxies at high-$z$ were much smaller than local galaxies with comparable mass (Daddi et al.~\\citeyear{daddi05}; Trujillo et al.~\\citeyear{trujillo06},~\\citeyear{trujillo07}; Cimatti et al.~\\citeyear{cimatti08}; van~Dokkum et al.~\\citeyear{vandokkum08}; Akiyama et al.~\\citeyear{akiyama08}; Franx et al.~\\citeyear{franx08}; Szomoru et al.~\\citeyear{szomoru10},~\\citeyear{szomoru12}; van~der~Wel et al.~\\citeyear{vanderwel11}; Barro et al.~\\citeyear{barro13}). At a fixed stellar mass, ETGs are claimed to have been significantly compact at high-$z$ and have evolved with rapid increase of their effective radius by a factor of $\\sim$ 4 or even larger from $z\\sim 2$ (Buitrago et al.~\\citeyear{buitrago08}; Carrasco et al.~\\citeyear{carrasco10}) and by a factor $\\sim$ 2 from $z \\sim 1$ (van~der~Wel et al.~\\citeyear{vanderwel08}; Trujillo et al.~\\citeyear{trujillo11}). To reach the size of local ETGs, rapid and violent evolutions by major merger (Hopkins et al.~\\citeyear{hopkins09a}) or minor merger (Bezanson et al.~\\citeyear{bezanson09}; Naab et al.~\\citeyear{naab09}) have been demanded.  Recent very deep infrared observations of high spatial resolution with the $Hubble\\ Space\\ Telescope\\ (HST)$ have shed light on morphological details and shapes of galaxies at high-$z$. Bruce et al.~(\\citeyear{bruce12}) studied over 200 massive galaxies at $1 < z < 3$ in the CANDELS-UDS field and found that these galaxies had much smaller size at a given mass than that of local ETGs. On the other hand, it has also been argued that the compact galaxies have apparent smaller effective radii because of low signal to noise ratio (S/N) (e.g., Ryan et al.~\\citeyear{ryan12}). The lack of the consideration for AGN component would also make the radius smaller (Yoshino \\& Ichikawa~\\citeyear{yoshino08}; Pierce et al.~\\citeyear{pierce10}). In addition, the best-fit morphological outputs with, for example, {\\ttfamily GALFIT} (Peng et al.~\\citeyear{peng02}), which is one of the most frequently used fitting codes for galaxy morphology, could be significantly changed with small differences of fitting inputs (e.g., initial guess, point spread function (PSF), weight image) and the image properties (e.g., size of postage stamp, sky background noise). Although {\\ttfamily GALFIT} are frequently used, it sometimes gives inappropriate results, mostly when used without careful considerations to image quality of galaxies and to the contamination by neighboring objects (H\\\"{a}u{\\ss}ler et al.~\\citeyear{haussler07}, hereafter H07; Barden et al.~\\citeyear{barden12}).  Some previous studies (e.g., Trujillo et al.~\\citeyear{trujillo06}; H07; Carollo et al.~\\citeyear{carollo13}; Mosleh et al.~\\citeyear{mosleh13}) estimated the errors in effective radius, $r_\\mathrm{e}$, and S\\'ersic index, $n$, using artificial galaxies (AGs), and derived simple relations between the original and output values. Szomoru et al.~(\\citeyear{szomoru10}) contrived to compensate the faint extended wings of galaxies. They estimated the limit of surface brightness and fitted the S\\'ersic profile to the galaxy images above the surface brightness limit with {\\ttfamily GALFIT}. Then, they corrected the result $r_\\mathrm{e}$ by calculating the residual counts between the original and model images. Based on a careful study of the bias of image quality and PSF profiles, van~der~Wel et al.~(\\citeyear{vanderwel12}) presented global structural parameters of more than 100,000 galaxies in the CANDELS survey. Bruce et al.~(\\citeyear{bruce12}) applied {\\ttfamily GALFIT} to three-component fitting (bulge, disk, and central components) for high-$z$ galaxies with a careful attention to the background noise and PSF convolution.  In addition, we should take account of the different analysis for local galaxies when comparing the morphological properties at high-$z$. The half-light radii of the SDSS local galaxies used in Shen et al.~(\\citeyear{shen03}, hereafter S03) were claimed to be underestimated (Guo et al.~\\citeyear{guo09}, hereafter G09; Simard et al.~\\citeyear{simard11}). The comparison of the size-stellar mass relations between the different definitions of stellar mass would also lead to inappropriate results (Mosleh et al.~\\citeyear{mosleh13}). The comparison of the structural parameters for high-$z$ galaxies with those in the local universe should be based on the consistent definition and analysis of galaxy data.  In this paper, we investigate the reliability and limit of {\\ttfamily GALFIT} to obtain the morphological properties of high-$z$ galaxies. Then, using deep near infrared (NIR) observations with Wide Field Camera 3 (WFC3) instrument installed on $HST$, we apply {\\ttfamily GALFIT} to galaxies in the Great Observatories Origins Deep Surveys-North (GOODS-N) region, in which Ichikawa et al.\\ (\\citeyear{ichikawa12}) (hereafter Ic12) studied the size evolution of galaxies in a non-parametric way with $K_\\mathrm{s}$-band ground-based images of MOIRCS Deep Survey (MODS). The ground-based images were not reliable enough for the morphological study of galaxies at $z>1$ with {\\ttfamily GALFIT} (Konishi et al.~\\citeyear{konishi11}). As such, Ic12 obtained the size-stellar mass relations based on half- and 90 percent light radii. On the other hand, deep images by WFC3 with much higher spatial resolution will allow us to apply {\\ttfamily GALFIT} to high-$z$ galaxies, including compact galaxies, for the morphological study.  An outline of the paper is as follows. In Section~\\ref{sec:sec2}, we describe the samples of massive galaxies in the GOODS-N. Using the background noise and PSF of WFC3 images, we make AGs with various shape parameters. We analyze them with {\\ttfamily GALFIT} and compare the results with the original parameters in Section~\\ref{sec:sec3}. We examine the reliability and the systematic errors of $r_\\mathrm{e}$ and $n$ of the AGs obtained with {\\ttfamily GALFIT} under some conditions. After examining the validity of the results, we apply the fitting method to massive galaxies in the GOODS-N. The results are described in Section~\\ref{sec:sec4} and Section~\\ref{sec:sec5}. Finally we discuss our results in comparison with those of previous parametric and non-parametric studies in Section~\\ref{sec:sec6}. Throughout this paper, we assume $\\Omega_m$ = 0.3, $\\Omega_\\mathrm{\\Lambda}$ = 0.7 and $H_0$ = 70 kms$^{-1}$Mpc$^{-1}$. We use the AB magnitude system (Oke \\& Gunn~\\citeyear{oke83}; Fukugita et al.~\\citeyear{fukugita96}). ",
        "conclusions": "\\label{sec:sec6}  We have obtained the size and shape of the massive galaxies, using MODS and $HST$/WFC3 CANDELS data in the GOODS-N region. Thanks to the high image quality and depth, we were allowed to analyze as faint objects as with $H_\\mathrm{AUTO}\\leq25$, which reaches less massive ($\\sim10^{10}$~M$_{\\sun}$) SFGs at $z\\sim2.5$ (QGs at $z\\sim1.5$) and massive ($\\geq10^{10.5}$~M$_{\\sun}$) SFGs at $z\\sim3.0$ (QGs at $z\\sim2.5$). With a careful test of {\\ttfamily GALFIT} analysis, applying magnitude criteria to AGs, we obtained unbiased morphological results for the samples. The tests for different PSFs and color effect also improved the reliability.  First, we discuss the evolution of $r_\\mathrm{e}$. As shown in Fig.~\\ref{fig:fig.6} and Fig.~\\ref{fig:fig.8}, \\textcolor{black}{ we found a number of compact QGs ($r_\\mathrm{e}\\sim1$~kpc or less) at $1.5 < z \\leq 2.5$, which is consistent with the previous studies (van~Dokkum et al.~\\citeyear{vandokkum09}; Szomoru et al.~\\citeyear{szomoru12}). } The size for massive QGs from $z\\sim2.5$ to $\\sim0.5$ is represented as $r_\\mathrm{e}\\propto (1+z)^{-\\alpha_{r_\\mathrm{e}}}$ with $\\alpha_{r_\\mathrm{e}} \\sim1.06$ (or a factor of $\\sim2.5$ increase from $z\\sim2.5$ to $\\sim0.5$ at a given stellar mass), which is consistent with the previous results of $\\alpha_{r_\\mathrm{e}} \\sim$ 0.7--1.5 (a factor of $\\sim$ 1.8--3.6 size increase). It is noted that SFGs have weaker size evolution ($\\alpha_{r_\\mathrm{e}} \\sim0.5$), irrespective of their stellar mass bins, which is inconsistent with previous studies (e.g., Mosleh et al.~\\citeyear{mosleh11}). However, the sample definition of SFGs in the studies are different ($UV$-bright galaxies in Mosleh et al.), and their results could not be compared with the present results. The scenario of the size evolution for QGs is in dispute for a decade both in theoretical (e.g., Fan et al.~\\citeyear{fan08}; Hopkins et al.~\\citeyear{hopkins09a}; Naab et al.~\\citeyear{naab09}; Oser et al.~\\citeyear{oser12}) and observational studies (Newman et al.~\\citeyear{newman10}; van~Dokkum et al.~\\citeyear{vandokkum10}). According to Eq.~4 in Naab et al.~(\\citeyear{naab09}), it takes $\\sim4~(\\sim2)$ minor merger events with 1:10 (3:10) mass ratio to explain the size evolution of the massive QGs (see also Bezanson et al.~\\citeyear{bezanson09}). In a merger scenario, however, about 10\\% of the massive galaxies at $z\\sim2$ are expected to have survived without \\textcolor{black}{ equal-mass merging (Hopkins et al.~\\citeyear{hopkins09a}), which leaves the superdense QGs morphologically unchanged to the local universe. } In the SDSS and other surveys for the local galaxies, however, there is not enough number of such a superdense relic (S03; Taylor et al.~\\citeyear{taylor10}; M\\'armol-Queralt\\'o et al.~\\citeyear{marmol12}). These facts suggest that the observed size evolution is hard to be explained only by the merger scenario (see also van~der~Wel et al.~\\citeyear{vanderwel09}; Barro et al.~\\citeyear{barro13}).  Secondly, we discuss the shapes of the galaxies by using $n$. As the shape of galaxy is believed to exhibit the evolution trace, irrespective of the galaxy masses, we discuss the evolution of $n$. In Fig.~\\ref{fig:fig.10}, we see that the typical $n$ of QGs is larger than that of SFGs. There is also difference between the evolution of $n$ of massive and less massive QGs, though the latter are limited to $z\\leq1.5$ due to the completeness limit. We found that $n$ of massive QGs significantly evolved with $\\alpha_n\\sim0.74$ of $n\\propto (1+z)^{-\\alpha_n}$, while those of less massive QGs and SFGs were unchanged on average ($\\alpha_n\\sim0.17$ and $\\sim 0$, respectively). The findings of the evolution in $n$ would give us a clue to understanding the morphological evolution of high-$z$ galaxies. \\textcolor{black}{ One of the most favorable scenarios is the inside-out mass growth, where the mass accretion in the outer parts of compact galaxies increase $r_\\mathrm{e}$ and $n$ toward low-$z$ (van~Dokkum et al.~\\citeyear{vandokkum10}). In van~Dokkum et al., they showed the evolution of $n$ ($n\\sim4$ to $\\sim2$ from $z\\sim2$ to $\\sim0.5$), as well as the size evolution ($r_\\mathrm{e}\\sim3$~kpc to 8~kpc over the same redshift range). } In addition to the observational results, numerical simulation by Naab et al.~(\\citeyear{naab09}) explained the evolution of $n$ by minor merging of satellite galaxies. As done in Chevance et al.~(\\citeyear{chevance12}), the investigation for the distribution and evolution in of $b/a$ over the redshift range may help our understandings for the low-$n$ QGs, which could be the progenitor of local bulges (see also Trujillo, Carrasco \\& Ferr\\'e-Mateu~\\citeyear{trujillo12}), though the small sample in the present study would not give the robust conclusion.  Based on the results of $r_\\mathrm{e}$ and $n$ in the present study, we then discuss the formation and evolution of massive QGs, taking account of the evolution of less massive QGs and SFGs as well. First, focusing on $n$, less massive SFGs at the whole redshift range ($0.5\\leq z\\leq2.5$) are thought to be the progenitor of the less massive QGs after exhausting their gas and passively evolve with little change in size and shape. In the context, the progenitors of comparatively less massive QGs at high-$z$ could be small or amorphous SFGs located at higher redshift ($z\\gtrsim3.0$) (Kajisawa \\& Yamada~\\citeyear{kajisawa01}), most of which are not included in the present study due to our detection limit. \\textcolor{black}{ While the sizes of QGs are smaller than those of SFGs at the whole redshift, it is believed that there are shrinkage of extended SFGs when they transform into QGs (Barro et al.~\\citeyear{barro13b}). The shrinkage of SFGs is also investigated by Dekel et al.~(\\citeyear{dekel13}), in which migration of star-forming clumps forms the blue compact galaxies (or blue nuggets), then quench into red nuggets (see also Noguchi~\\citeyear{noguchi99}). We need more detailed investigation whether the compact galaxies keep $n$ lower after the migration of star-forming clumps, while Williams et al.~(\\citeyear{williams14}) showed that blue nuggets had $\\langle n\\rangle \\sim$ 2-3 and red nuggets $\\langle n\\rangle \\sim$ 3-4. Whitaker et al.~(\\citeyear{whitaker12}) found very little difference in the sizes of young and old quenched galaxies. }  After the birth of low-$n$ red QGs, in the later epoch ($z<1.5$) major or minor mergers of those dry (gas-poor) QGs dominates the evolution of massive QGs, enlarging their size rapidly over the cosmic time (Gao et al.~\\citeyear{gao04}). The dry merger is also consistent with the evolution of $n$ because it is efficiently change the light profile of galaxy (Barnes~\\citeyear{barnes92}; Hernquist~\\citeyear{hernquist92}; Naab et al.~\\citeyear{naab09}), while wet (gas-rich) merger would reproduce disk-like galaxies (Steinmetz \\& Navarro~\\citeyear{steinmetz02}; Springel \\& Hernquist~\\citeyear{springel05}; Robertson et al.~\\citeyear{robertson06}; Cox et al.~\\citeyear{cox06}). On the other hand, we found no significant evolution in $r_\\mathrm{e}$ and $n$ of less massive QGs. \\textcolor{black}{ To explain this, we might need to consider the environmental effect on their formation (or halo mass size; Dekel et al.~\\citeyear{dekel13}). However, we avoid discussing this because of the small sample and limited redshift range of the less massive QGs in the present study. }   \\textcolor{black}{ It should be noted that the argument above (see also Carollo et al.~\\citeyear{carollo13}) is complicated by the fact that samples of massive galaxies, not separated into star-forming and quiescent, have been demonstrated to grow smoothly with time (e.g., van~Dokkum et al.~\\citeyear{vandokkum10}; Patel et al.~\\citeyear{patel13}). To corroborate the evolution scenario avobe, the discussion based on comoving number density for each group of galaxies would be helpful. Patel et al.~(\\citeyear{patel13}) studied the structural evolution of massive galaxies by linking progenitors and descendants at a constant cumulative number density. However, small number of samples in the present study would not be enough for the robust discussion in statistical sense. The field variance in a small field would also hamper the reliable conclusion. In addition, the present study investigated the evolution of $r_\\mathrm{e}$ at given two mass bins while those previous studies investigated the evolution taking account of the stellar mass evolution ({\\it differential size evolution}) by using the constant number density method or mass-normalized size. The mass-normalized size is useful not only to see the differential size evolution, but also to compensate the weak completeness, which could give rise to the biased (or spurious) size evolution. In Fig.~\\ref{fig:norm} we show the mass-normalized size evolution for QGs and SFGs. $r_\\mathrm{e,norm}$ is derived as $r_\\mathrm{e,norm}=r_\\mathrm{e}/(M_*^\\mathrm{cor}/M_\\mathrm{c})^{\\alpha_\\mathrm{M}}$, where ${\\alpha_\\mathrm{M}}$ is the linear slope of the size-stellar mass relation for each redshift bin (see Section~\\ref{sec:sec5.1}). It is noted that the relations are derived only using the sample within the completeness limits. Then, we see the size evolution with $r_\\mathrm{e}\\propto(1+z)^{-\\alpha_{r_\\mathrm{e}}}$ finding that $\\alpha_{r_\\mathrm{e}}=1.07\\pm0.15$ for QGs and $\\alpha_{r_\\mathrm{e}}=0.59\\pm0.05$ for SFGs, which is consistent with the results of massive QGs and SFGs at fixed mass bins. The results ensure our results of the size evolution over the redshift ranges, while we need more sample to investigate the weak evolution of $r_\\mathrm{e}$ (and $n$) of less massive QGs. }  \\begin{figure} \\figurenum{12} \\plotone{norm.eps} \\caption{ Mass-normalized size evolution for QGs (top) and SFGs (bottom). Black dotted lines represent the regressions of $r_\\mathrm{e,norm}\\propto(1+z)^{-\\alpha_{r\\mathrm{e}}}$. } \\label{fig:norm} \\end{figure}   Finally, we should refer to the scatters in size at whole redshift range comparable to that of the local populations (Fig.~\\ref{fig:fig.9}). This suggests that there had already been massive QGs at $z\\sim$ 1.5--3, whose morphological properties were similar to that of the local ones. This is encouraged by the fact that the velocity dispersion of QGs at $z\\sim1.6$ were found to be comparable to those of the local galaxies (Onodera et al.~\\citeyear{onodera12}). It should be noted, however, that a single S\\'ersic profile may not be appropriate to fit galaxies at high-$z$, even if they seem to be well-virialized. Bruce et al.~(\\citeyear{bruce12}) applied {\\ttfamily GALFIT} with two components (bulge, disk) massive galaxies ($M_*>10^{11}$~M$_{\\sun}$) at $1<z<3$ and found many bulge$+$disk systems. In the local universe, Huang et al.~(\\citeyear{huang13}) found the elliptical galaxies were well fitted by three components. However, multi-component fitting for our sample is too complicated because we of the small apparent size of, in particular, galaxies with $M_*\\sim10^{10}$~M$_{\\odot}$ at $z>1$, making it hard to find the \"best-fit\" results, while we adopted a single component S\\'ersic fit in the present study. \\textcolor{black}{ To discuss the origin of the scatters in size, we need further investigation with, for example, mass-normalized size, which is beyond the present study. }"
    },
    "1402/1402.1470.txt": {
        "abstract": "{We present a {\\it Herschel} far-IR and sub-mm study of a sample of 120 galaxies in 28 Hickson compact groups. Fitting their UV to sub-mm spectral energy distributions with the model of da Cunha et al. (2008), we accurately estimate the dust masses, luminosities, and temperatures of the individual galaxies. We find that nearly half of the late-type galaxies in dynamically ``old'' groups, those with more than 25\\% of early-type members and redder UV-optical colours, also have significantly lower dust-to-stellar mass ratios compared to those of actively star-forming galaxies of the same mass found both in HCGs and in the field. Examining their dust-to-gas mass ratios, we conclude that dust was stripped out of these systems as a result of the gravitational and hydrodynamic interactions, experienced owing to previous encounters with other group members. About 40\\% of the early-type galaxies (mostly lenticulars), in dynamically ``old'' groups, display dust properties similar to those of the UV-optical red late-type galaxies. Given their stellar masses, star formation rates, and UV-optical colours, we suggest that red late-type and dusty lenticular galaxies represent transition populations between blue star-forming disk galaxies and quiescent early-type ellipticals. On the other hand, both the complete absence of any correlation between the dust and stellar masses of the dusty ellipticals and their enhanced star formation activity, suggest the increase in their gas and dust content due to accretion and merging. Our deep {\\it Herschel} observations also allow us to detect the presence of diffuse cold intragroup dust in 4 HCGs. We also find that the fraction of 250$\\mu$m emission that is located outside of the main bodies of both the red late-type galaxies and the dusty lenticulars, is 15-20\\% of their integrated emission at this band. All these findings are consistent with an evolutionary scenario in which gas dissipation, shocks, and turbulence, in addition to tidal interactions, shape the evolution of galaxies in compact groups. } ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%% Since most galaxies are found within large structures, such as groups and clusters, studying how the environment can affect their properties (i.e. stellar populations and morphology) is crucial to understanding their evolution. It is now known that dynamical interactions and the merging of galaxies can affect both their morphologies (induced bars, bridges, tails, and other tidal distortions) and the star formation and nuclear activity \\citep[i.e.][]{Struck99}. Nevertheless, galaxy environments can differ depending on the number and the kinematic behaviour of their members. Compact groups of galaxies are systems of several galaxies that display high galaxy densities, similar to those found in the central regions of rich clusters, unlike clusters, they have much lower average velocity dispersions (\\citealt{Hickson97} quotes a $\\sigma \\sim$250 km/s). There are exceptions, such as Stephan's Quintet, where an intruder is colliding with the group with a velocity difference of over 800 km/s \\citep{Hickson92}. As a result, galaxies in compact groups can experience a series of strong tidal encounters with other group members during their lifetimes. For these reasons, compact groups are considered ideal systems for studying the effects of dense environments in galaxy evolution. To examine the properties of galaxies in compact groups we used the catalogue defined by \\citet{Hickson82}; it consists of 100 groups, containing 451 galaxies, in compact configurations (less than 5 arcmin), within relatively isolated regions where no excess of other surrounding galaxies can be seen. Even though all Hickson compact groups (HCGs) were selected as having four or more galaxies, the original sample was later reduced to 92 groups, since spectroscopic observations revealed the inclusion of interlopers among their members \\citep[see][]{Hickson92}.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.50]{figure_example.png} \\caption{Spitzer/IRAC ``true'' colour image of HCG 07 (panel a). The blue channel traces the 3.6$\\mu$m emission, the green the 4.5$\\mu$m, and the red the 8.0$\\mu$m. The {\\it Herschel}/PACS 160$\\mu$m (panel b), and the {\\it Herschel}/SPIRE 250$\\mu$m (panel c) images of the same group. In the bottom left corners of panels b and c we include the beam sizes of the {\\it Herschel} instruments (11$''$ and 22$''$, respectively). In panel c) we also display for comparison the beams of the {\\it AKARI} 160$\\mu$m band (70$''$) and {\\it IRAS} 100$\\mu$m (1.2$'\\times$5$'$). } \\label{fig:fig_example} \\end{center} \\end{figure*}  Detailed studies in the past three decades have revealed that HCGs occupy a unique position in the framework of galaxy evolution. They display an excess of elliptical galaxies (relative to the field), with their spiral galaxy fraction nearly half of what it is observed in the field (43\\%; \\citealt{Hickson82}). \\citet{Mendes94} showed that almost half of their galaxies display morphological features of interactions, and \\citet{Zepf93} revealed indications of mergers in the irregular isophotes of the elliptical galaxies. In addition, \\citet{Ponman96} found hot X-ray gas in $\\sim$75\\% of the groups, implying that their members reside in a common dark matter halo, while \\citet{Desjardins13} show that the X-ray luminosity increases as group HI-to-dynamical mass ratio decreases. However, the hot gas is not in hydrostatic equilibrium and, thus, these systems are not the low-mass analogues of rich groups or clusters. Single-dish radio measurements reveal that galaxies in compact groups are generally deficient in HI with a median mass two times less than what is observed in loose groups \\citep{Verdes01}. Furthermore, \\citet{MartinezBadenes12} show an increase in the conversion of HI-to-H$_{2}$ caused by the on-going tidal interactions in the spiral galaxies of groups found in early stages of evolution, followed by HI stripping and the decrease in molecular gas because of the lack of replenishment.  Nearly 40\\% of HCG members for which nuclear spectroscopy has been obtained display evidence of active galactic nuclei \\citep{Shimada00, Martinez10}.  Although tidal stripping is clearly a strong contender for explaining HI deficiencies, there are cases where the gas seems to have been almost completely stripped from the galaxies, most likely involving both hydrodynamic, as well as pure gravitational, processes. The best studied example is Stephan's Quintet (SQ = HCG92), where HI and molecular studies show that almost all of the cold ISM now resides in the intragroup medium \\citep{Williams02,Gao00,Guillard12}. Strong dissipative processes, in addition to gravity, are clearly at work on a large scale in this system, leading to extensive shock heating in the interfaces between galaxies and the intragroup medium, as is evident by the presence of hot X-ray gas, warm molecular hydrogen, and diffuse ionized carbon \\citep{Trinchieri05,O'Sullivan09,Appleton06,Cluver10,Appleton13}.  While systems like SQ seem rare, they are not unique and may represent a rapid hydrodynamic phase in which head-on collisions between the group members can cause rapid sweeping of the gas (e. g. \\citealt{Appleton96} for the ring galaxy VII Zw 466) in addition to disruptive tides, conventional mergers, and gas depletion. For example, a recent study of HCG57 by Alatalo et al. (in preparation) also shows a highly disruptive collision between two galaxies in which the molecular and atomic gas is dramatically changed by direct collision. This may be similar to the Taffy system \\citep{Condon94}, where a head-on collision has pulled gas into a highly turbulent bridge between the galaxies, heating the molecular gas \\citep{Peterson12}.  \\citet{Cluver13} studied the mid-IR spectra of 74 galaxies in 23 HCGs and discovered that more than 10\\% of them contain enhanced warm molecular hydrogen emission that, like SQ, may be indicative of shocks and turbulent heating of their ISM. Furthermore, these H$_{2}$ enhanced galaxies (called molecular hydrogen emission galaxies -- MOHEGs by \\citealt{Ogle10}) seem to primarily occupy a special place in mid-IR colour space referred to as the ``gap'' by \\citet{Johnson07}.  Although not strictly a ``gap'' \\citep[see][]{Bitsakis11},  Johnson et al. as well as \\citet{Walker12} argue that there is a deficiency of galaxies in the {\\it Spitzer} IRAC colour-colour space defined primarily as lying in the range $-0.4<log(f_{5.8}/f_{3.6}) < 0.0$  which they attribute to rapid evolution from the star forming to the quiescent galaxy sequences. In addition, \\citet{Walker13} showed that galaxies located in the area of the mid-IR gap, have already transitioned to the optical {\\it [g-r]} red sequence. Nevertheless, since these authors do not correct these colours for dust attenuation, some of the galaxies can be found in the red sequence due to extinction. \\citet{Cluver13} name the region as the ``infrared green valley'' since it falls between the two main peaks in the mid-IR colour space characterised primarily by star-forming galaxies and galaxies dominated by old stellar populations in compact groups, and as this current paper shows, galaxies in this region tend to be a mix of  dusty red late-type and dusty lenticular galaxies. The discovery that the majority of the galaxies that lie in this region of colour space are also MOHEGs, may lend some credence to the idea that galaxies in the green valley are not just the tail of two separate galaxy populations, but may be galaxies ``in transition'' from an evolutionary perspective. These authors suggest that these shock-excited systems may be heated by collisions with HI clouds liberated into the intragroup medium by previous tidal stripping. SQ could be seen as an extreme example of this process. However, as we see later in this paper, re-accretion of gas by early-type galaxies could also cause evolution of galaxies in the other direction into the green valley.  In \\citet{Bitsakis10,Bitsakis11} -- hereafter $paper~I$ and $paper~II$, respectively -- we obtained the mid-IR fluxes of a sample of 32 HCGs (containing 135 galaxies) to study, for first time, their star formation properties. Using all the multi-wavelength data available, combined with the physically-motivated model of \\citet{daCunha08}, we were able to consistently interpret the UV-to-mid-IR spectral energy distribution (SED) of the galaxies in terms of a number of physical parameters, such as their bolometric luminosity, star formation rate, stellar mass, and dust attenuation. We also identified the dynamical state of each group based on the fraction of late-type galaxies they contain, defining as ``old'' (or ``young'') those with more (or less) than 75\\% early-type members, and showed that this is consistent with classification based on HI gas content and morphology. We concluded that spiral galaxies in dynamically young groups do not differentiate their star formation properties and UV-optical colours from similar galaxies in other environments, such as in the field and in early stage interacting pairs. On the other hand, the dynamically old late-type galaxies display lower specific star formation rates (sSFR) and redder {\\it [NUV-r]} colours than what we see in the field. We suggested that this can be attributed to past galaxy-galaxy interactions and possible minor merging of the group members since the time of its formation, which has increased their stellar masses. Nevertheless, the idea of red late-type galaxies populating the ``green valley'' is not new. \\citet{Boselli08} show that ram pressure stripping can remove substantial amounts of atomic gas out of the low-luminosity (L$_{H}<$10$^{9.6}$L$_{H,\\odot}$) late-type galaxies entering the Virgo cluster, quenching their star formation and making them display redder colours relatively short timescales ($<$150Myr). Later, this work has been expanded by \\citet{Hughes09}, \\citet{Cortese09}, and \\citet{Gavazzi10} to brighter galaxies in high density environments (i.e. in the clusters of Virgo, Coma, and A1367), showing that the red late-type galaxies display high HI deficiencies and truncation of star formation. Despite the difference in the mechanisms that cause this ``reddening'' (ram pressure in clusters and gravitational stripping in groups) it seems that the effects on stellar populations and star formation activity are similar in compact environments.  In our earlier work the estimates of the far-IR luminosities were based on limited information available at the time for generating SEDs in the far-IR. Our knowledge of the far-IR was based primarily on 60 and 100$\\mu$m measurements from IRAS (which had an $\\sim$4 arcmin beam shape), making it difficult to properly separate galaxies in compact groups that often also lie well within the beam profile. As a result, much of our far-IR knowledge of individual galaxies was necessarily based on extrapolation from the Spitzer 24$\\mu$m data points.  To address this problem, we use in the current paper data taken from several {\\it Herschel} Open Time proposals led by our team, as well as other archival data sets from the {\\it Herschel} Science Archive\\footnote {The Herschel Science Archive is publicly available at http://herschel.esac.esa.int/Science\\_Archive.shtml } to make the first detailed analysis of 28 HCG, extending the SED to the far-IR and sub-mm wavelengths. In Section 2, we describe the selection of the samples, used in this work, the data reduction processes, and the UV to far-IR SED modelling of the galaxies. In Section 3, we present the results of our analysis. In Section 4, we discuss  the implications of these findings, and in Section 5 we present our conclusions.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In this paper we have presented our results of the first far-IR and sub-mm analysis of Hickson compact group galaxies, based on {\\it Herschel} observations. We fitted the UV-to-far-IR SEDs and accurately estimated for first time the mass and temperature of the dust for a sample of 120 galaxies, contained in 28 HCGs, as well as a control sample of field galaxies, the KINGFISH sample. Based on the analysis of these results we conclude the following:  \\begin{itemize}  \\item Late-type galaxies that were found to display red {\\it [NUV-r]} colours, owing to the enhanced stellar populations they contain, also have significantly lower dust-to-stellar mass ratios than normal HCG or field ``blue cloud'' late-type galaxies of the same stellar mass. Examining them in more detail, we find that their atomic gas-to-dust mass ratios are more than an order of magnitude lower than the blue cloud late-type galaxies, even though they have similar molecular gas-to-dust mass ratios. We suggest that significant amounts of dust were removed from the disks of these galaxies as a result of tidal and hydrodynamic interactions.  \\item Unlike to the dynamically ``young'' groups, where no early-type galaxy member was detected, almost 40\\% of the early-type members of the dynamically ``old'' groups were detected in all the far-IR and sub-mm bands of {\\it Herschel}. The ``dusty'' lenticular galaxies display very similar dust masses to the red late-type galaxies. Based on the stellar masses, the star formation rates, the {\\it [NUV-r]} and mid-IR colours, and the dust-to-stellar and gas mass ratios of the red late-types and the dusty lenticulars, we suggest that they represent transition populations between blue star-forming disk galaxies and quiescent early-type ellipticals in groups. Moreover, MOHEGs seem to dominate this phase of transition.  \\item We detect seven ``dusty'' elliptical galaxies, which display high dust-to-stellar mass ratios, comparable to or higher than those of the red late-type galaxies. Unlike to the rest of the sample, there is no correlation between their dust and stellar masses. Moreover, these galaxies display significantly higher sSFRs than the remaining ellipticals in HCGs. We suggest that in these enhanced systems, their dust content is due to accretion from other group members and the intragroup medium, as well as to active star formation.  \\item Examining the spatial distribution of the emission in the red late-type galaxies, we find that $\\sim$20\\% of their 250$\\mu$m emission is located outside their main bodies. Similarly, nearly half of the dusty lenticulars also display such extended emission, with the remaining having dust disks that are much smaller than their stellar ones. Estimating the mass of the extended dust we find that it is not enough to compensate for the very low dust-to-stellar mass ratios of these galaxies.  \\item Searching for the missing dust, we conclusively detect cold intragroup dust in at least four HCGs. Depending on the assumptions we make for its temperature and emissivity, we estimate its mass to be approximately 7-76\\% of the total dust mass of the group. Comparing the atomic gas and the dust maps and examining the detection limits, we also conclude that there is no correlation in the distributions of the intragroup gas and the cold dust.  \\end{itemize}  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.75]{figure12.png} \\caption{True colour images, created using the 3.6, 4.5, and 8.0\u00b5m {\\it Spitzer}/IRAC mosaics of groups HCG16, HCG54, HCG79, and HCG92, respectively. The green and red contour maps are the 250$\\mu$m {\\it Herschel}/SPIRE and the {\\it VLA}/HI moment-0 maps of the same groups. The five contours correspond to 3, 5, 7, 9, and 20$\\sigma$. In the bottom right corner of each panel, we present the size of the SPIRE 250$\\mu$m and {\\it VLA} band beams, respectively. The R.A. and Dec. scales of the axes are in J2000.0 coordinates.} \\label{fig:fig12} \\end{center} \\end{figure*}"
    },
    "1402/1402.2728_arXiv.txt": {
        "abstract": "With improvements in exoplanet detection techniques, the number of multiple planet systems discovered is increasing, while the detection of potentially habitable Earth-mass planets remains complicated and thus requires new search strategies. Dynamical studies of known multiple planet systems are therefore a vital tool in the search for stable and habitable planet candidates.  Here, we present a dynamical study of the three-planet system HD~204313 to determine whether it could harbour an Earth-like planet within its habitable zone for a sufficient time to develop life. We found two semi-stable regions in the system, but neither prove stable for long enough for a terrestrial planet to develop life.  Our investigations suggest that overlapping weak and high order resonances may be responsible for these semi-stable regions. This study established a framework for a larger project that will study the dynamical stability of the habitable zone of multiple planet systems, providing a list of interesting targets for future habitable low-mass planet searches. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5614_arXiv.txt": {
        "abstract": "{% Bubble N107 was discovered in the infrared emission of dust in the Galactic Plane observed by the Spitzer Space Telescope (GLIMPSE survey: $l \\approx 51\\fdg0$, $b \\approx 0\\fdg1$). The bubble represents an example of shell-like structures found all over the Milky Way Galaxy. }{% We aim to analyse the atomic and molecular components of N107, as well as its radio continuum emission. With the help of numerical simulations, we aim to estimate the bubble age and other parameters which cannot be derived directly from observations. }{% From the observations of the \\ion{H}{i} (I-GALFA) and \\element[][13]{CO} (GRS) lines we derive the bubble's kinematical distance and masses of the atomic and molecular components. With the algorithm DENDROFIND, we decompose molecular material into individual clumps. From the continuum observations at $1420\\ \\mathrm{MHz}$ (VGPS) and $327\\ \\mathrm{MHz}$ (WSRT), we derive the radio flux density and the spectral index. With the numerical code \\emph{ring}, we simulate the evolution of stellar-blown bubbles similar to N107. }{% The total \\ion{H}{i} mass associated with N107 is $5.4 \\times 10^3\\ M_\\odot$. The total mass of the molecular component (a mixture of cold gasses of H$_2$, CO, He and heavier elements) is $1.3 \\times 10^5\\ M_\\odot$, from which $4.0 \\times 10^4\\ M_\\odot$ is found along the bubble border. We identified 49 molecular clumps distributed along the bubble border, with the slope of the clump mass function of $-1.1$. The spectral index of $-0.30$ of a strong radio source located apparently within the bubble indicates nonthermal emission, hence part of the flux likely originates in a supernova remnant, not yet catalogued. The numerical simulations suggest N107 is likely less than $2.25\\ \\mathrm{Myr}$ old. Since first supernovae explode only after $3\\ \\mathrm{Myr}$ or later, no supernova remnant should be present within the bubble. It may be explained if there is a supernova remnant in the direction towards the bubble, however not associated with it. }{% } ",
        "introduction": "Bubble N107 (SIMBAD: \\object{CPA2006 N107}) is one of the largest bubbles in the catalogue by \\citet{2006apj...649..759c}. This catalogue is based on the infrared observations made with the Spitzer Space Telescope, (GLIMPSE survey, \\citealp{2003pasp..115..953b}) and contains more than 300 bubbles found in the interstellar medium (ISM) near the Galactic Plane. The catalogue was later expanded by \\citet{2012mnras.424.2442s} to more than 5000 bubbles, which were identified by a community of over $35\\,000$ volunteers. These bubbles, with typical radii less than $10\\ \\mathrm{pc}$, are examples of shell-like structures -- features found in the ISM all over our Galaxy. Other examples are \\ion{H}{i} shells in catalogues by \\citet{1979apj...229..533h,1984apjs...55..585h,2005a&a...437..101e}, with sizes up to $3500\\ \\mathrm{pc}$; or dusty ``loops'' discovered in the IRAS full-sky survey by \\citet{2004a&a...418..131k} and \\citet{2007a&a...463.1227k}.  The varying sizes and masses of these structures suggest different mechanisms responsible for their creation. These mechanisms, which do not have to act solely, include:  \\begin{enumerate} \\item feedback from massive (OB) stars: radiation, winds and supernovae (see refs.\\ below); \\item infall of high-velocity clouds into the Galactic disc (e.g. \\citealp{1984apjs...55..585h,1988ara&a..26..145t,1996a&a...313..478e}); \\item energy and mass inserted by gamma-ray bursts (e.g. \\citealp{1998apj...501l.163e,1998apj...503l..35l,1999a&a...350..457e}); and \\item turbulence (\\citealp{2005apj...630..238d}). \\end{enumerate}  \\noindent Bubble N107, as most of the GLIMPSE bubbles, is likely a result of stellar feedback \\citep{2006apj...649..759c,2010a&a...523a...6d}, which may follow three different models:  \\begin{enumerate} \\item An expanding \\ion{H}{ii} region, assuming radiative heating and ionisation of hydrogen by UV photons from a massive star \\citep{1954ban....12..187k,1955apj...121....6O,1978ppim.book.....s}. The pressure inside the \\ion{H}{ii} region is higher compared to the ambient neutral medium, so the \\ion{H}{ii} region expands and collects the neutral material in an expanding shell. \\item A wind-blown bubble, supposing that energy is released via stellar winds \\citep{1975apj...200l.107c,1977apj...218..377w}. Such a bubble, formed around a massive star, can be divided into several layers: an innermost layer, where the wind is freely expanding up to a layer of shocked stellar wind, where its kinetic energy is thermalised.  This layer of the shocked wind is followed by a layer of swept up ISM, which in later stages of evolution is mainly atomic and molecular. \\item A supernova explosion, supposing an abrupt energy input into the ISM, producing an expanding shell/remnant \\citep[][and citations therein]{1974apj...188..501c}. \\end{enumerate}  \\noindent These different models of stellar feedback were also discussed by \\citet{1988ara&a..26..145t,1996a&a...313..478e} and \\citet{1999a&a...350..457e}.  The shell-like structures can trigger star formation, since they consist of a layer of cold and dense material. In the so-called collect-and-collapse scenario \\citep{1977apj...214..725e,1994apj...427..384e,1994mnras.268..291w}, the shell fragments and creates a new generation of stars. Radiation driven implosion \\citep{1990a&a...233..190d,1994a&a...289..559l} considers the compression of preexisting condensations (globules) by the pressure of the ionised gas. A brief review of these and other triggering mechanisms is given in a study of GLIMPSE bubbles by \\citet{2010a&a...523a...6d}.   In this paper, we present a multiwavelength study of one of the largest GLIMPSE bubbles: N107, which lies in the central plane of the Galactic disc ($l \\approx 51\\fdg0$, $b \\approx 0\\fdg1$). A multiwavelength, colour picture of N107 is shown in fig.~\\ref{fig_n107_multi_co}.  The paper is organised in the following way: In section~\\ref{sec_observation}, we complement the Spitzer infrared observations with \\ion{H}{i} line observations from the I-GALFA survey \\citep{2010hia....15..788k}, \\element[][13]{CO} ($J=1\\mathrm{-}0$) line observations from the GRS survey \\citep{2006apjs..163..145j} and radio continuum observations at $1420$ and $327\\ \\mathrm{MHz}$ from the VGPS \\citep{2006aj....132.1158s} and WSRT \\citep{1996apjs..107..239t} surveys. From the line observations, we derive the bubble's LSR\\footnote{LSR = local standard of rest} radial velocity (RV), kinematical distance and masses of the molecular and atomic components. From the radio continuum observations, we derive the flux densities and corresponding spectral indices of two radio sources lying apparently inside the bubble. In section~\\ref{sec_analysis_of_molecular_clumps}, we use a recently published algorithm DENDROFIND \\citep{2012a&a...539a.116w} to decompose the molecular gas associated with the bubble into individual clumps and derive the slope of the clump mass function. In section~\\ref{sec_simulations}, we use the numerical code \\emph{ring} \\citep{1990iaus..144p.101p,1996a&a...313..478e} to simulate the bubble's evolution in order to estimate its age, energy input and the size and mass of the molecular cloud from which the bubble evolved. Finally, a discussion and conclusions are presented in sections~\\ref{sec_discussion} and \\ref{sec_conclusion}.   \\begin{figure*} \\includegraphics[width=\\hsize]{fig_n107_multi_co} \\caption{Bubble N107 -- false colour multiwavelength image.  Red corresponds to $24\\ \\mathrm{\\upmu m}$ continuum, green to $8\\ \\mathrm{\\upmu m}$ continuum and blue to \\element[][13]{CO} ($J=1\\mathrm{-}0$) line integrated over the radial velocities of $38.5$ to $47.6\\ \\mathrm{km/s}$. Note that the $24\\ \\mathrm{\\upmu m}$ continuum (red) is dominated by the emission of very small dust grains heated by a nearby massive star. The $8\\ \\mathrm{\\upmu m}$ continuum (green) is dominated by the emission of PAHs -- large molecules found within molecular clouds. The emission of PAHs is tracing the photodissociation regions (PDRs) -- edges of these clouds illuminated by UV radiation from a massive star. The CO line (blue) traces the cold gas of molecular clouds. The distinct red-green structure seen in the lower-left in the direction of the bubble's opening is likely not related to the bubble, since the radial velocity of the associated molecular material ($\\approx 60\\ \\mathrm{km/s}$) differs significantly from the radial velocity of the N107 complex ($\\approx 43\\ \\mathrm{km/s}$). } \\label{fig_n107_multi_co} \\end{figure*} ",
        "conclusions": "}  We studied dust bubble N107, discovered by \\citet{2006apj...649..759c} with infrared observations by the Spitzer Space Telescope. Radio observations of the \\ion{H}{i} and \\element[][13]{CO} ($J=1\\mathrm{-}0$) lines revealed its atomic and molecular components, which allowed us to determine the bubble radial velocity, kinematical distance and mass of these components. Using the code DENDROFIND \\citep{2012a&a...539a.116w} we decomposed the molecular ring associated with the bubble into 49 individual clumps and found that the clump mass spectrum slope is about $-1.1$. Radio continuum observations at $1420$ and $327\\ \\mathrm{MHz}$ suggest the radio flux coming from the bubble direction has both thermal and nonthermal components; i.e., besides a classical \\ion{H}{ii} region, a supernova remnant contributes to the emission. We simulated the evolution of stellar-blown bubbles within molecular clouds with the numerical code \\emph{ring} \\citep{1990iaus..144p.101p,1996a&a...313..478e}. We were able to produce two groups of bubbles with a similar angular mass distribution and expansion velocity as N107: \\textbf{(A)} Bubbles formed within an oblate molecular cloud ($\\sigma_z/\\sigma_{xy} = 6/12$), with the energy input $\\dot E < 0.4 \\times 10^{50}\\ \\mathrm{erg/Myr}$ and the evolution time $t \\approx 1.75 \\dots 2.25\\ \\mathrm{Myr}$. \\textbf{(B)} Bubbles formed within a prolate molecular cloud ($\\sigma_z/\\sigma_{xy} = (20 \\dots 28)/12$), with the energy input $\\dot E \\approx 2.7 \\times 10^{50}\\ \\mathrm{erg/Myr}$ and the evolution time  $t = 1\\ \\mathrm{Myr}$. Considering the estimated bubble evolution time ($< 2.25\\ \\mathrm{Myr}$), no supernova remnant should be present inside the bubble, yet. This may be explained by a supernova remnant present in the direction towards the bubble, however not associated with it. A summary of the parameters we derived for bubble N107, together with the parameters given by \\citet{2006apj...649..759c}, is presented in table~\\ref{tab_summary}.   \\begin{table} \\centering \\caption{ Summary of parameters derived for N107. The first section gives the values from the original catalogue of \\citet{2006apj...649..759c}. The second section gives the values derived from observation (sec.~\\ref{sec_observation}). The third section gives the values estimated from simulations (sec.~\\ref{sec_simulations}). } \\label{tab_summary}  \\begin{tabular}{ll} \\hline\\hline Property  &  Value \\\\ \\hline \\\\ \\hline \\multicolumn{2}{l}{From catalogue of \\citet{2006apj...649..759c}} \\\\ \\hline Galactic longitude (central)   & $50\\fdg972$ \\\\ Galactic latitude (central)    & $0\\fdg078$ \\\\ Angular radius (mean)          & $11\\farcm39$ \\\\ Angular thickness (mean)       & $2\\farcm26$ \\\\ \\\\ \\hline \\multicolumn{2}{l}{Derived from observations} \\\\ \\hline Radial velocity (central)        & $43\\ \\mathrm{km/s}$ \\\\ Expansion velocity               & $\\ 8\\ \\mathrm{km/s}$ \\\\ Distance (kinematical)           & $3.6\\ \\mathrm{kpc}$ \\\\ Radius (mean)                    & $11.9\\ \\mathrm{pc}$ \\\\ Mass in \\ion{H}{i} (aperture 1)  & $(5.4 \\pm 1.0) \\times 10^3\\ M_\\odot$ \\\\ \\multicolumn{2}{l}{Mass of molecular component (aperture 1)} \\\\ & $(1.3 \\pm 0.4) \\times 10^5\\ M_\\odot$ \\\\ \\multicolumn{2}{l}{Mass of molecular component (aperture 2)\\tablefootmark{1}} \\\\ & $(4.0 \\pm 1.2) \\times 10^4\\ M_\\odot$ \\\\ \\\\ \\hline \\multicolumn{2}{l}{Estimated from simulations} \\\\ \\hline Energy input rate              & $< 0.4 \\times 10^{50}\\ \\mathrm{erg/Myr}$\\tablefootmark{A} \\\\ & $\\approx 2.7 \\times 10^{50}\\ \\mathrm{erg/Myr}$\\tablefootmark{B} \\\\ Dislocation of expansion centre & $\\approx 16 \\dots 22\\ \\mathrm{pc}$\\tablefootmark{A} \\\\ & $\\approx 10 \\dots 12\\ \\mathrm{pc}$\\tablefootmark{B} \\\\ Parental cloud thickness in LOS  & $\\approx 6\\ \\mathrm{pc}$\\tablefootmark{A} \\\\ & $\\approx 20 \\dots 28\\ \\mathrm{pc}$\\tablefootmark{B} \\\\ Parental cloud total mass      & $\\approx 5.8 \\dots 15.5 \\times 10^5\\ M_\\odot$ \\\\ Evolution time (age)           & $\\approx 1.75 \\dots 2.25\\ \\mathrm{Myr}$\\tablefootmark{A} \\\\ & $\\approx 1 \\mathrm{Myr}$\\tablefootmark{B} \\\\ \\hline \\end{tabular}  \\tablefoot{ \\tablefoottext{1}{Aperture 2 covers only the immediate vicinity of the bubble. See fig.~\\ref{fig_apertures} for an illustration.} \\tablefoottext{A}{Estimated from simulations from group A.} \\tablefoottext{B}{Estimated from simulations from group B.} }  \\end{table}"
    },
    "1402/1402.3518_arXiv.txt": {
        "abstract": "{Searching for disk-outflow systems in massive star-forming regions is a key study to assess the main physical processes in the recipe of massive star formation.} {We have selected a hot molecular core (HMC) in the high-mass star-forming region \\G23, where VLBI multi-epoch observations of water and methanol masers have suggested the existence of rotation and expansion within 2000~AU from its center. Our purpose is to image the thermal line and continuum emission at millimeter wavelengths to establish the physical parameters and velocity field of the gas in the region.} {We performed SMA observations at 1.3~mm with both the most extended and compact array configurations, providing sub-arcsecond and high sensitivity maps of various molecular lines, including both hot-core and outflow tracers. We also reconstruct the spectral energy distribution of the region from millimeter to near infrared wavelengths, using the Herschel/Hi-GAL maps, as well as archival data.} {From the spectral energy distribution, we derive a bolometric luminosity of $\\sim$$\\rm4\\times10^4~L_\\odot$. Our interferometric observations reveal that the distribution of dense gas and dust in the HMC is significantly flattened and extends up to a radius of 8000~AU from the center of radio continuum and maser emission in the region. The equatorial plane of this HMC is strictly perpendicular to the elongation of the collimated bipolar outflow, as imaged on scales of $\\sim$0.1--0.5~pc in the main CO isotopomers as well as in the SiO(5--4) line. In the innermost HMC regions ($\\rm \\la1000~AU$), the velocity field traced by the CH$_3$CN\\,($\\rm 12_K-11_K$) line emission shows that molecular gas is both expanding along the outflow direction following a Hubble-law, and rotating about the outflow axis, in agreement with the (3-D) velocity field traced by methanol masers. The velocity field associated with rotation indicates a dynamical mass of $\\sim$$\\rm 19~M_{\\odot}$ at the center of the core. The latter is likely to be concentrated in a single O9.5 ZAMS star, consistent with the estimated bolometric luminosity of \\G23. The physical properties of the CO\\,(2--1) outflow emission, such as its momentum rate $6\\times10^{-3}$~M$_\\odot$~km~s$^{-1}$~yr$^{-1}$ and its outflow rate $2\\times10^{-4}$~M$_\\odot$~yr$^{-1}$, support our estimates of the luminosity (and mass) of the embedded young stellar object.} {} ",
        "introduction": "The formation process of high-mass stars ($\\rm \\ga 10~M_{\\odot}$) is still to be established. While theoretical models have been proposed, only recently substantial progress has been  made observationally (see, e.g., \\citealt{Zinnecker2007}). It is widely accepted that hot molecular cores (HMCs) are the sites where massive stars form and it is thus of interest to study these objects at high angular resolution, through high-density tracers such as rare molecular species. In the recipe of star formation a crucial role is played by disks and the associated jets/outflows, which in the case of high-mass stars might reduce the radiation pressure barrier on the infalling gas (i.e., the so-called flashlight effect) thus making accretion onto the protostar possible (see, e.g., \\citealt{Kuiper2013}).  For these reasons, many observations have been devoted to HMCs and their molecular environment (e.g., \\citealt{Cesaroni2007} and references therein). The success of this type of studies depends mostly on the selection of appropriate targets, as various facts may hinder the analysis of the velocity field inside the HMCs. For example, too large a distance makes the linear resolution insufficient to identify any putative disk in the HMC, while a suitable inclination with respect to the line of sight is needed to detect the velocity gradient due to the outflow expansion and/or disk rotation. Multi-epoch VLBI observations of maser lines provide us with 3-D information on the velocity field on scales as small as 10--100~AU and may be used to select promising objects for disk/outflow studies. With this in mind, we have performed interferometric observations of dense gas and outflow tracers towards the high-mass star-forming region (HMSFR) \\G23. This star forming site harbors a known HMC that we have recently studied by means of high-resolution maser observations, suggesting the presence of rotation and expansion within 2000~AU from the HMC center \\citep{Sanna2010,Sanna2012}.   \\G23\\ is a HMSFR located at a trigonometric distance of 4.59$^{+0.38}_{-0.33}$~kpc \\citep{Brunthaler2009} which exceeds a bolometric luminosity of $\\rm 10^4~L_\\odot$ (see \\citealt{Sanna2010} for a recent review).  On a parsec scale, Spitzer/GLIMPSE images of the field reveal bright, ``green fuzzy'' emission at 4.5~$\\mu$m, believed to be associated with an intense outflow activity (e.g., \\citealt{Cyganowski2009}). Indeed, a massive $^{12}$CO outflow has been detected with the IRAM interferometer at sub-pc scales (Furuya et al.\\,2008). Towards the center of the molecular outflow, emission from high energy transitions of molecular species such as CH$_3$CN and NH$_3$ has been detected, indicating the existence of a HMC hosting newly formed massive stars (\\citealt{Codella1997,Furuya2008}). All the strongest maser species known to date have been detected from the HMC and subsequently imaged with multi-epoch VLBI observations \\citep{Sanna2010,Sanna2012}. We hence decided to perform Submillimeter Array (SMA) observations with sub-arcsec resolution to unveil the gas structure and kinematics in the HMC and associated outflow on scales ranging from 0.01 to 0.1~pc, with the aim to compare these results to those obtained from the maser study.  This paper is organized as follows. In Sect.~2, we give the details of our SMA observations. In Sect.~3, we present the results of the analysis of selected molecular lines for typical HMC and outflow tracers.  We also collect the photometric data at different IR wavelengthes at the position of \\G23, from several Galactic surveys available to date, and reconstruct its spectral energy distribution. In Sect.~4, we discuss the structure and velocity field of the region from the HMC to the outflow scale and the nature of the young stellar object(s) (YSOs) embedded in the HMC. Finally, the main results are summarized in Sect.~5. ",
        "conclusions": "We used the SMA to image various hot-core and outflow tracers from the high-mass star forming region \\G23\\ at 1.3~mm. We have also combined the Herschel/Hi-GAL data with ancillary archival data to reconstruct the SED associated with the HMC in \\G23\\ (between~3.4~$\\mu$m and~1.1~mm) and derive its luminosity. In this paper, we have focused on a comparison of the gas dynamics in the inner HMC, on scales $\\rm \\la 0.01~pc$, with respect to that inferred from typical outflow tracers, on scales $\\rm \\ga 0.1~pc$. Our main results can be summarized as follows:  \\begin{enumerate} \\item We find that the HMC core is significantly flattened (width-to-length ratio $\\sim2$), which corresponds to an inclination of $30\\degr$ under the assumption that we are observing an oblate core tilted with respect to the line of sight. Also, the size of the core appears to decrease with excitation energy of the observed transition, suggesting the presence of a temperature gradient due to embedded YSOs. From the mean CH$_3$CN spectrum within a radius of $\\sim$3000~AU from the HMC center, we infer a rotational temperature of $\\sim$195~K. \\item The CH$_3$CN emission appears to trace two distinct velocity fields, one consistent with expansion along the outflow axis, the other suggesting rotation about the same axis (possibly Keplerian), in good agreement with the (3-D) velocity field traced by methanol masers. The central mass implied by the rotational component is $\\sim$19~$M_\\odot$, after taking into account an inclination of the rotation axis of $30\\degr$. This value is consistent with the bolometric luminosity of the \\G23\\ region ($\\sim$$4\\times10^4~L_\\odot$) estimated by fitting the SED. \\item We have estimated the Tommre $Q$ stability parameter at two different radii in the putative disk, and find that $Q$ decreases significantly from small to large radii, consistent with the disk being more stable close to the central star(s), as expected. \\item The outflow emission imaged in the CO isotopomers and SiO(5--4) line well matches the direction of the rotation axis (P.A.=\\,$58\\degr$) of the flattened HMC. The outflow is well collimated with a width-to-length ratio for each lobe of about 0.3. The energetics of the molecular outflow, such as its momentum rate of $6\\times10^{-3}$~M$_\\odot$~km~s$^{-1}$~yr$^{-1}$, is in good agreement with that expected for the measured bolometric luminosity.   \\end{enumerate}  We conclude that the HMC in \\G23\\ is likely hosting an O9.5 (proto)star lying at its center, powering the bipolar outflow observed on the parsec scale, and shows clear proof of rotation up to a radius of $\\sim$1000~AU. Higher angular resolution observations are needed in order to confirm that the HMC is indeed hosting a single massive star, to trace the inner rotation curve, and to find out whether accretion is still going on inside the HMC."
    },
    "1402/1402.1492_arXiv.txt": {
        "abstract": "Radio relics are diffuse radio sources observed in galaxy clusters, probably produced by shock acceleration during cluster-cluster mergers. Their large size, of the order of 1 Mpc, indicates that the emitting electrons need to be (re)accelerated locally. The usually invoked Diffusive Shock Acceleration models have been challenged by recent observations and theory. We report the discovery of complex radio emission in the Galaxy cluster PLCKG287.0 +32.9, which hosts two relics, a radio halo, and several radio filamentary emission. Optical observations suggest that the cluster is elongated, likely  along an intergalactic filament, and displays a significant amount of substructure. The peculiar features of this radio relic are that (i) it appears to be connected to the lobes of a radio galaxy and (ii) the radio spectrum steepens on either side of the radio relic. We discuss the origins of these features in the context of particle re-acceleration. ",
        "introduction": "The hierarchical model of structure formation predicts that clusters of galaxies grow by successive accretion of smaller sub-units. A large amount of energy is dissipated in the intra-cluster medium (ICM) as a result of merger events, and a fraction of this energy could be channeled into the amplification of the magnetic fields \\citep[see e.g.][for a review]{Dolag08} and into the acceleration of high-energy Cosmic Rays \\citep[see][for a review]{Petrosian08} that in turn may produce observable extended radio emission. Extended radio sources on a cluster scale, not associated with any optical counterpart, have been detected in a number of galaxy clusters \\citep{Venturi07, Giovannini09,2011A&A...533A..35V,Bonafede12,Feretti12}. They are called radio halos and radio relics, depending on their morphology, location and radio properties.\\\\  \\noindent {\\bf Radio relics}\\\\  Radio relics are extended sources, characterised by a steep spectrum\\footnote{The spectrum is here defined as $S(\\nu) \\propto \\nu^{-\\alpha}$.} ($\\alpha >$1) and strong polarisation ($\\sim$ 20-30 \\% at 1.4 GHz). Their origin is subject to debate and not yet understood. There is a general consensus that they are related to shock waves, occurring in the ICM during mergers. Shock waves should be able to amplify magnetic fields and accelerate electrons up to relativistic energies, hence producing synchrotron radio emission  \\citep{Bruggen11,2012MNRAS.423.2781I,Vazza12,sk13}. Although this picture is roughly consistent with observations, some important issues remain unexplained. First, a radio relic is not always detected when a shock wave is present in the ICM, as revealed by X-ray observations \\citep{Russell11} and recently \\citet{2013MNRAS.433..812O} have found a displacement between the X-ray shock wave and the radio relic in 1RXS J0603.3+4214. Secondly, most of the relics do not appear to be co-located with a shock wave in X-ray observations (see review by \\citealt{Bruggen11} and ref. therein). Thirdly, shock waves in the ICM are characterised by Mach numbers of the order of 2 - 4 \\citep{sk08,va09shocks} so that their efficiency in accelerating particles is expected to be too low to account for the radio emission.\\\\ Some authors have proposed that shock waves by themselves are not sufficient to accelerate the particles from the thermal pool to relativistic energies. This would indicate that a seed population of relativistic electrons must be already present before the shock passage \\citep{KangRyu11,KangRyu12,Pinzke13}. This seed population of old radio plasma could be re-energised by a shock wave, explaining why the connection between a relic and a shock wave is not one to one. Since the old radio plasma is much more energetic than the thermal gas in the ICM, a low acceleration efficiency could be sufficient to re-accelerate the particles and power the radio emission. The problem is now to find the source for this seed population of relativistic electrons. Some authors have suggested that they come from a previous episode of shock acceleration (\\citealt{Macario11}). Another possibility is that the old plasma comes from the activity of radio galaxies.\\\\  Recent results by \\citet{VazzaBruggen13} have shown that the gamma-ray upper limits to the cluster  emission from the {\\it Fermi} satellite are at odds with the hypothesis that radio relics are generated by Diffusive Shock Acceleration (DSA) mechanism, at least in the way it appears to work for supernova remnants. In fact, if shock waves accelerate electrons in the ICM, also protons should be accelerated, with a higher efficiency. In inelastic collisions with the thermal protons, relativistic protons produce  pions and gamma-ray photons, which {\\it Fermi} should have detected.\\\\  \\noindent {\\bf Radio Halos}\\\\  Radio halos are diffuse Mpc-sized objects found in merging galaxy clusters \\citep[e.g.][]{Buote01,Cassano10} and characterised by steep spectra (see e.g. reviews by \\citealt{Ferrari08,Feretti12}). Two different classes of models have been proposed so far: the hadronic models \\citep{2010ApJ...722..737K,2011A&A...527A..99E} and the turbulent re-acceleration models \\citep{Fujita03,CassanoBrunetti05}. In this paper, we study the peculiar and complex  emission of the cluster PLCKG287.0 +32.9 to test  DSA models. PLCKG287.0 +32.9  is a massive cluster ($M_{500} \\sim 1.4 \\times 10^{15}$ solar masses, Planck Collaboration 2013) located at redshift $z=0.39$, showing a disturbed X-ray morphology (Bagchi et al. 2011) which indicates that the cluster is undergoing a merger event. The X-ray centre of the cluster has a Right  Ascension equal to $11h50m51.02s$ and a declination of  $-28d04'09.37\"$ PLCKG287.0 +32.9 is known to host two radio relics  \\citep{Bagchi11}. Our new radio observations give a different picture of the system. The radio emission is more complex and more extended than previously observed, and offers a unique opportunity to unravel the origins of radio relics. The paper is organised as follows: In Sec.  \\ref{radio} we present our new observations. The X-ray  and optical data are discussed in Sec. \\ref{sec:xray} and \\ref{sec:optical}, respectively. In Sec. \\ref{radioanalysis} the radio properties are studied, and the spectral properties are presented in Sec. \\ref{sec:spix}. Results are discussed in \\ref{discussion} and we present our conclusions in Sec. \\ref{conclusions}.\\\\ Throughout this paper, we assume a concordance $\\rm{\\Lambda CDM}$ cosmological model, with $H_0=$ 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=$ 0.27, and $\\Omega_{\\Lambda}=$ 0.73. One arcmin corresponds to 316 kpc at $z=$0.39.    \\begin{figure}  \\vspace{90pt} \\begin{picture}(90,90)   \\put(-20,0){\\includegraphics[width=9cm]{otticoradiox.eps}}     % \\end{picture} \\caption{X-ray emission in red (XMM-Newton), radio emission at 323 MHz in blue (low resolution, beam FWHM $\\sim 22'' \\times 18''$) and in green (high resolution:  beam FWHM$\\sim 13'' \\times 8''$).} \\label{fig:xray} \\end{figure}          \\begin{table*} \\centering \\caption{Radio observations.} \\begin{tabular}{c c c c c c c c} \\hline Frequency & Radio telescope & Observing date &  time  & bandwidth  & Int. time & rms noise & Beam \\\\ MHz        &                &                         &    h        &    MHz          &    s                 & mJy/beam &\\\\ \\hline 325   &  GMRT    &    Jan 2013 & \t 7        & 33             &    8 \t\t  & 0.1               & 13$'' \\times 8''$    \\\\ 610   &  GMRT    &    May 2013&         10     & 33              &    8\t\t   & 0.065          & 7$'' \\times 5''$ \\\\ 150    &  GMRT    &\tJul  2011 &        8     & 16              &    8\t\t   & 1.3               &25$'' \\times 18''$  \\\\ 3000   &  JVLA      &\tJun 2013 &         6\t  & 2000         &    3                 & 0.05    & 19$'' \\times 11''$  \\\\ \\hline  \\multicolumn{7}{l}{\\scriptsize  Col. 1: Observing  Frequency, Col. 2: Radio telescope. Col. 3: Date of the observation. Col. 4: total observing time. }\\\\ \\multicolumn{7}{l}{\\scriptsize  Col. 5: Observing bandwidth. Col 6: Integration time per visibility.}\\\\ \\multicolumn{7}{l}{\\scriptsize Col 7: 1$\\sigma$ rms noise reached in the high resolution images.}\\\\ \\multicolumn{7}{l}{\\scriptsize Col 8: FWHM of major and minor axes of the restoring beam in the high resolution images. }\\\\ \\end{tabular} \\label{tab:obs} \\end{table*} ",
        "conclusions": "\\label{conclusions} We have presented a multi-wavelength analysis of the galaxy cluster PLCKG287.0 +32.9, that might shed new light on the origin of radio relics. Our results can be summarised as follows: \\begin{itemize} \\item{Optical data suggests that PLCKG287.0 +32.9 is located within a  6-Mpc long intergalactic filament. The cluster has undergone a major merger (NWc -SEc), slightly misaligned with respect to the main direction of accretion. Along the filament two sub-clumps are detected. The galaxies along the intergalactic filament, as well as the two sub-clumps (SEext and NW),  are likely on the process of accretion onto the main cluster.} \\item{Two radio relics and a radio halo are present in the cluster. Additional emission is detected NW of the relic. Another relic is located SE of the cluster, at a projected distance of  2.8 Mpc. A radio halo is located in the cluster center, and filamentary emission is detected around both relics.} \\item{The large projected distance of the SE relic and the small projected distance on the NW relic are interpreted as shock waves produces at different times during a minor merger event. A sub-cluster (with a mass $\\sim$ 0.1 of the main one) could have caused a first shock wave at his first infall onto the main cluster, and a second shock wave during the second core-passage (NW relic). The radio halo could be due to the major merger between the SEc and NEc sub-clusters.} \\item{The NW relic emission fades into the lobes of a radio galaxy, indicating that radio relics originate from electrons previously injected by AGN and reaccelerated, likely by a shock wave. The spectral index analysis supports this interpretation.  However, it is necessary to assume that two or more radio sources have injected electrons into the pre-shock region to account for the relic size.} \\end{itemize}"
    },
    "1402/1402.1171_arXiv.txt": {
        "abstract": "We present a generalized Newtonian description of particle dynamics valid for any spherically symmetric, static black hole spacetime. This approach is derived from the geodesic motion of test particles in the low-energy limit. It reproduces exactly the location of the marginally stable, marginally bound, and photon circular orbits; the radial dependence of the energy and angular momentum of circular orbits; parabolic motion; pericentre shift; and the spatial projection of general trajectories. As explicit examples of the new prescription, we apply it to the Schwarzschild, Schwarzschild--de Sitter, Reissner--Nordstr{\\\"o}m, Ay{\\'o}n-Beato--Garc{\\'i}a, and Kehagias--Sfetsos spacetimes. In all of these examples, the orbital and epicyclic frequencies are reproduced to better than $10\\%$.  The resulting equations of motion can be implemented easily and efficiently within existing Newtonian frameworks. ",
        "introduction": "Since its birth almost 100 years ago, the theory of general relativity (GR) has successfully withstood continuous experimental testing: from the anomalous perihelion precession of Mercury and the bending of light rays from distant stars by the Sun to the indirect confirmation of the existence of gravitational waves from the orbital decay of the Hulse--Taylor pulsar \\citep{will06}. All of these examples, however, only probe the weak field limit of the theory while the strong field regime remains largely uncharted \\citep{psaltis}. One of the most dramatic consequences of the strong field limit of GR is the prediction of black holes and, although current astrophysical observations are consistent with their existence, the available evidence does not allow to discriminate between the objects predicted by GR and those coming from alternative gravitational theories \\citep[see, e.g.][]{barausse13,bambi13}.  It is generally believed that astrophysical black holes possess a certain degree of intrinsic angular momentum, either since birth or as a result of subsequent accretion processes \\citep{narayan}. For this reason, it is expected that the spacetime around astrophysical black holes should be well approximated by the solution of a rotating black hole found by \\cite{kerr}. Nevertheless, many key relativistic features can already be investigated in the non-rotating, spherically symmetric black hole solution of Schwarzschild. Moreover, one of the first steps in the development of an alternative theory of gravity is to look for the kind of black hole solutions that it admits, and, given their simplicity, a prominent role in this search is played by static, spherically symmetric black hole spacetimes.  The Schwarzschild black hole solution of GR is the best known example of a static and spherically symmetric spacetime, but there exist other metrics that might be relevant for different astrophysical scenarios, especially in the context of alternative theories of gravity and/or extra degrees of freedom for the central object. Examples include \\begin{itemize} \\item[{\\bf-}] The Reissner--Nordstr\\\"om solution describing a charged black hole within Einstein--Maxwell equations. For a discussion of particle motion in this spacetime see, e.g~\\cite{bicak89,pugliese,grunau}. \\item[{\\bf-}] The Schwarzschild--de Sitter spacetime describing a Schwarzschild black hole in an expanding universe with cosmological constant. See \\cite{stuchlik08,hackmann08b,hackmann08a} for a general discussion of the motion of test particles in this spacetime. \\item[{\\bf-}] The Kehagias--Sfetsos spacetime \\citep{kehagias}, a black hole solution of the Lorentz-violating gravitational theory of \\cite{horava}. The motion of test particles in this spacetime has been discussed in, e.g.~\\cite{abdu,enolskii,vieira}. \\item[{\\bf-}] The Boulware--Deser black hole solution in (4+1)-dimensional Gauss--Bonnet gravity \\citep{boulware}. \\item[{\\bf-}] Regular black holes. These are curvature singularity-free, exact solutions of Einstein equations coupled to some nonlinear electrodynamics \\citep[see, e.g.][]{bardeen68,ABG}. See \\cite{zhou,garcia} for a discussion of test particle motion in these spacetimes. \\item[{\\bf-}] A Schwarzschild black hole pierced by a cosmic string \\citep{aryal}. For a discussion of geodesic motion in this spacetime see, e.g.~\\cite{hackmann}. \\item[{\\bf-}] Schwarzschild and Reissner--Nordstr\\\"om black holes in spacetimes with higher dimensions \\citep{tangherlini,emparan}. For a general discussion of geodesic motion in these spacetimes see, e.g.~\\cite{hackmann08}. \\end{itemize}  In this article we present a generalized Newtonian description of the motion of test particles in any given static, spherically symmetric spacetime. This work extends the scheme of \\cite{TR} (referred to as Paper~I in the following) where, by considering the low-energy limit of geodesic motion, we introduced an  accurate Newtonian description of particle motion around a Schwarzschild black hole. The approach presented in Paper~I reproduces exactly certain key relativistic features of this spacetime (e.g.~radial location, angular momentum and energy of distinct circular orbits), while several other properties (e.g.~Keplerian and epicyclic frequencies) are described with a better accuracy than with commonly used pseudo-Newtonian potentials. Moreover, it showed a very good agreement with the thin disc accretion model of \\cite{novikov73} and the analytic model for relativistic accretion of \\cite{tejeda2,tejeda3}.  The paper is organized as follows. In Section~\\ref{S2} we present the model and discuss the motion of general test particles within the new approach. In Section~\\ref{S3} we derive explicit expression for the equations of motion that can be easily implemented within existing Newtonian frameworks. Circular orbits are discussed in detail in Section~\\ref{S4} while perturbations away from them are considered in Section~\\ref{S5}. Finally, we summarize our results in Section~\\ref{S6}. ",
        "conclusions": "\\label{S6}  We have presented a generalized Newtonian description of the motion of test particles valid for any static, spherically symmetric spacetime based on the low-energy limit of geodesic motion. This work constitutes a generalization of the approach developed in Paper~I for a Schwarzschild black hole. Among the various examples of spherically symmetric spacetimes that might be of relevance for different astrophysical contexts, we have considered the use of the new approach for a Schwarzschild black hole with cosmological constant (Schwarzschild--de Sitter spacetime), a charged black hole in Einstein--Maxwell theory (Reissner--Nordstr\\\"om spacetime), a charged regular black hole in Einstein theory coupled with a non-linear electrodynamics (Ay\\'on-Beato--Garc\\'ia spacetime), and a black hole solution of the Lorentz-violating gravitational theory of Ho{\\v r}ava-Lifshitz (Kehagias-Sfetsos spacetime).  We have shown that the new approach reproduces exactly the radial location of the photon, marginally bound, and marginally stable circular orbits. Moreover, the radial dependence of the binding energy and angular momentum of circular orbits is also reproduced exactly. In addition, the spatial projection and the pericentre shift of general particle trajectories coincide exactly with the corresponding relativistic expressions. On the other hand, we have shown that, for the spacetimes considered in this paper, the Keplerian and epicyclic frequencies are reproduced to better than $10\\%$ for all of the radii that allow stable circular motion.  We believe that the ability of the present approach to reproduce all of these key relativistic features simultaneously and, in addition, its ease of implementation within existing Newtonian frameworks, make it a simple, yet powerful tool for studying astrophysical accretion flows onto a broad family of black hole spacetimes.  \\ack  We would like to thank John C.~Miller and Zdenek Stuchl\\'ik for insightful discussions and critical comments on the manuscript. Part of this work was started during the conference Prague Synergy 2013. ET would like to thank the organizers for their hospitality and the participants for stimulating discussions. This work has been supported by the Swedish Research Council (VR) under the grant 621-2012-4870.  \\begin{landscape} \\begin{table} \\caption{Application to examples of static and spherically symmetric black hole spacetimes. The function $\\alpha$ is connected with the spacetime metric as in \\eq{e1}. The ratio $\\Omega_*/\\Omega = \\Omega_*^\\parallel/\\Omega^\\parallel$ corresponds to the function $\\xi$ defined in \\eq{e22}. The second to last column gives the maximum percentage error obtained for $\\Omega_*$, $\\Omega_*^\\parallel$ and $\\Omega_*^\\perp$ (we have considered only values of $r$ for which circular motion is stable, see Figure~\\ref{f1}).} \\label{t1} \\begin{tabular}{llllp{5.7cm}} \\br\\\\[-7pt] Spacetime & $\\alpha$ & $\\frac{\\Omega_*}{\\Omega}$ & $\\frac{\\Omega-\\Omega_*}{\\Omega}$  & Notes \\\\[6pt]\\mr\\\\[-7pt] Schwarzschild & $1 - \\frac{ 2\\,\\rg }{ r }$ & $\\frac{r-2\\,\\rg}{\\sqrt{r(r-3\\,\\rg)}}$ & $\\leqslant 5.7\\%$ & $\\rg = \\G M/\\cc^2$ is the gravitational radius and $M$ is the mass of the black hole.\\\\ % Schwarzschild--de Sitter & $1 - \\frac{ 2\\,\\rg }{ r } - \\frac{\\Lambda\\,r^2}{3}$ & $\\frac{r-2\\,\\rg-\\Lambda\\,r^3/3}{\\sqrt{r(r-3\\,\\rg)}}$ & $\\leqslant 7.1\\%$ & $\\Lambda$ is the cosmological constant. The condition $\\Lambda\\,\\rg^2 \\leqslant 7.11\\times10^{-4}$ should be satisfied in order to have stable circular orbits. \\\\ % Reissner--Nordstr\\\"om & $ 1 - \\frac{ 2\\,\\rg }{ r } +\\frac{\\rQ^2}{r^2}$ & $\\frac{r^2-2\\,\\rg\\,r+\\rQ^2}{r\\sqrt{r^2-3\\,\\rg\\,r+2\\,\\rQ^2}}$ & $\\leqslant 8.2\\%$ & $\\rQ = \\sqrt{\\G}\\,Q/\\cc^2$ and $Q$ is the electrical charge of the black hole. To prevent the formation of a naked singularity, the condition $|\\rQ| < \\rg$ should be satisfied.\\\\ % Ay\\'on-Beato--Garc\\'ia & $ 1 -\\frac{ 2\\,\\rg \\,r^2}{\\left(r^2+\\rQ^2\\right)^{3/2} }+\\frac{ \\rQ^2r^2}{\\left(r^2+\\rQ^2\\right)^{2} }$ & $\\frac{1 -\\frac{ 2\\,\\rg \\,r^2}{\\left(r^2+\\rQ^2\\right)^{3/2} }+\\frac{ \\rQ^2r^2}{\\left(r^2+\\rQ^2\\right)^{2} }}{\\left[1 -\\frac{ 3\\,\\rg \\,r^4}{\\left(r^2+\\rQ^2\\right)^{5/2} }+\\frac{ 2\\,\\rQ^2r^4}{\\left(r^2+\\rQ^2\\right)^{3} }\\right]^{\\nicefrac{1}{2}}}$ & $\\leqslant 7.5\\%$ & We have considered values of $q\\leqslant q_\\s{\\text{crit}} \\simeq 0.634\\,\\rg$ in order to ensure a monotonic dependence of the ISCO on $q$. \\\\ % Kehagias-Sfetsos & $ 1 +\\omega\\,r^2\\left[ 1-\\left(1+\\frac{4\\,\\rg}{\\omega\\,r^3}\\right)^{\\nicefrac{1}{2}}\\right]$ & $\\frac{1 +\\omega\\,r^2\\left[ 1-\\left(1+\\frac{4\\,\\rg}{\\omega\\,r^3}\\right)^{\\nicefrac{1}{2}}\\right]}{\\left[1-\\frac{3\\,\\rg}{r}\\left(1+\\frac{4\\,\\rg}{\\omega\\,r^3}\\right)^{-\\nicefrac{1}{2}}\\right]^{\\nicefrac{1}{2}}}$ & $\\leqslant 6.3\\%$ & $\\omega$ gives the deviation away from GR (which is recovered in the limit \\mbox{$\\omega\\rightarrow \\infty$}). The condition $\\omega\\,\\rg^2\\geqslant 1/2$ should be satisfied to prevent the emergence of a naked singularity.\\\\ \\br \\end{tabular} \\end{table} \\end{landscape}"
    },
    "1402/1402.6798_arXiv.txt": {
        "abstract": "We explore the inflationary phase of a scalar field with a kinetic term non-minimally coupled to gravity. We find that one of the slow-roll conditions is naturally consequence of the equation of motion of the scalar field. Thus, slow-roll conditions impose fewer constraints on potentials than other inflationary models. Moreover, it is demonstrated that the inflationary phase can be described by just one slow-roll parameter. By investigating the metric perturbations, it is shown that except for one potential, almost all potentials have the same pattern in the ($n_{s}$, $r$) plane. We provide an exact solution for the exceptional case. The exact solution represents the condensed scalar field and results in an accelerated expansion. \\pacs{98.80.Cq} ",
        "introduction": "INTRODUCTION} Observation of the cosmic microwave background and large scale structure are consistent with slow-roll inflation paradigm in which a scalar field rolls slowly down its potential \\cite{plank-data}.\\\\ To have slow-roll inflationary phase, certain slow-roll conditions are considered that lead to some constraints on parameters of a model \\cite{weinberg}. For example, consider a simple theory of inflation described by the following Lagrangian \\begin{equation}\\label{simple} \\mathcal{L_{s}}=\\sqrt{-g}\\left[\\frac{R}{2\\kappa^{2}}-\\frac{1}{2}g^{\\mu\\nu}\\partial_{\\mu}\\phi\\partial_{\\nu}\\phi-V(\\phi)\\right], \\end{equation} where $\\kappa^{2}=8\\pi G$, in the flat Friedmann-Robertson-Walker metric with signature $(-,+,+,+)$ and $a$ as the scale factor. There exist two \\emph{independent} slow-roll conditions. The first of these conditions is actually the definition of inflationary phase states that $\\dot{H}/H^{2}\\equiv\\epsilon\\ll1$. Since the evolution of the scalar field is given by a second-order differential equation, it is not clear that the first condition holds over an extended period. So, the second condition is required in which we \\emph{demand} that $\\ddot{\\phi}\\ll H\\dot{\\phi}$.\\\\ Also, to determine the details of models from observations, various slow-roll parameters are used such as $\\epsilon, \\eta\\equiv\\dot{\\epsilon}/H\\epsilon, \\chi\\equiv\\dot{\\eta}/H\\eta $. The parameters are used to determined the specific potential term in \\eqref{simple} for the inflation era \\cite{plank-data}. Since the shape of the potential term is important to study high energy physics, the values of these parameters connect cosmology to particle physics.\\\\ In this paper we introduce a model for which we need just one parameter, $\\epsilon$, to describe slow-roll inflation phase. The model is given by the following action \\begin{equation}\\label{0-1} S=\\int d^{4}x\\sqrt{-g}\\left[\\frac{R}{2\\kappa^{2}}-g^{\\mu\\nu}\\partial_{\\mu}\\varphi\\partial_{\\nu}\\varphi++\\frac{1}{2}\\alpha^{2}G^{\\mu\\nu}\\partial_{\\mu}\\varphi\\partial_{\\nu}\\varphi-V(\\varphi) \\right], \\end{equation} where $G^{\\mu\\nu}$ is the Einstein's tensor and $\\alpha$ is an inverse mass parameter. We have chosen $+$ sign for the second term because for the de Sitter space we have $G_{\\mu\\nu}\\propto-g_{\\mu\\nu}$, so the scalar filed in \\eqref{0-1} has the same dynamics as the scalar filed in \\eqref{simple} in the de Sitter space.\\\\ The third term is one of the operators of the Horndeski's scalar-tensor theory \\cite{Horndeski}. Thus, in this model we confront with second-order differential equations. C. Germani et al studied $V(\\varphi)=\\lambda\\varphi^4$ for \"the new higgs inflation\" \\cite{higgs,higgse} and extended their model for other cases in \\cite{higgs-extend}. It has been shown that a new type of inflationary phase arises when the third term dominates over the standard kinetic term \\cite{higgs,higgs-extend}. So in this paper we will focus on this type of inflation and we will neglect the second term in the action \\eqref{0-1}.\\\\ Compared to the other works \\cite{higgs,higgs-extend}, we have two new results. The first result is that for $V(\\varphi)=M^{6}/\\varphi^2$ background equations can be solve exactly and metric perturbations will be obtained. The second result is to show that up to the first order in $\\epsilon$, in contrast to the simple model of inflation \\eqref{simple}, the values of measurable quantities of the inflationary regime of this model are almost independent from the form of the potential term in \\eqref{0-1}, as is shown in Fig. 1. Also, the model gives values for the parameters which are agreement with the Planck data in \\cite{plank-data}.\\\\ The organization of this paper is as follows: in Sec. II an exact solution for a special potential is given and the metric perturbations of the solution is discussed. we show that the scalar field is condensed and results in accelerated expansion. Sec. III is devoted to study the background and the metric perturbations of general potentials. We show that many features of the exact solution are similar to other potentials. We summarize our findings and discuss about the results in Sec. IV. In Appendix the second order actions for the metric perturbations, in which $V(\\varphi)\\neq0$, is given. ",
        "conclusions": ""
    },
    "1402/1402.6525.txt": {
        "abstract": "{In protoplanetary disks, the inner boundary between the turbulent and laminar regions could be a promising site for planet formation, thanks to the trapping of solids at the boundary itself or in vortices generated by the Rossby wave instability. %% Activity at this %% radius will also influence the thermodynamic structure of much of the %% dead-zone (including features such as the ice line). At the interface, the disk thermodynamics and the turbulent dynamics are entwined because of the importance of turbulent dissipation and thermal ionization. Numerical models of the boundary, however, have neglected the thermodynamics, and thus miss a part of the physics.} {The aim of this paper is to numerically investigate the interplay between thermodynamics and dynamics in the inner regions of protoplanetary disks by properly accounting for turbulent heating and the dependence of the resistivity on the local temperature.} {Using the Godunov code RAMSES, we performed a series of 3D global numerical simulations of protoplanetary disks in the cylindrical limit, including turbulent heating and a simple prescription for radiative cooling.} {We find that waves excited by the turbulence significantly heat the dead zone, and we subsequently provide a simple theoretical framework for estimating the wave heating and consequent temperature profile. In addition, our simulations reveal that the dead-zone inner edge can propagate outward into the dead zone, before staling at a critical radius that can be estimated from a mean-field model. The engine driving the propagation is in fact density wave heating close to the interface. A pressure maximum appears at the interface in all simulations, and we note the emergence of the Rossby wave instability in simulations with extended azimuth.} {Our simulations illustrate the complex interplay between thermodynamics and turbulent dynamics in the inner regions of protoplanetary disks. They also reveal how important activity at the dead-zone interface can be for the dead-zone thermodynamic structure.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.6090_arXiv.txt": {
        "abstract": "{We combined new near-infrared VLT/HAWK-I data of the globular clusters (GCs) in the isolated edge-on S0 galaxy \\object{NGC\\,3115} with optical and spectroscopic ones taken from the literature, with the aim of analyzing the multiband GC color distributions.  A recent study from the SLUGGS survey has shown that the GCs in this galaxy follow a bimodal distribution of Ca II triplet indices. Thus, \\object{NGC\\,3115} presents a critical example of a GC system with multiple, distinct, metallicity subpopulations, and this may argue against the ``projection'' scenario, which posits that the ubiquitous color bimodality mainly results from nonlinearities in the color-metallicity relations.  Using optical, NIR, and spectroscopic data, we found strong and consistent evidence of index bimodality, which independently confirms the metallicity bimodality in \\object{NGC\\,3115} GCs.  At the same time, we also found evidence for some color--color nonlinearity.  Taken in the broader context of previous studies, the multicolor consistency of the GC bimodality in \\object{NGC\\,3115} suggests that in cases where GC systems exhibit clear differences between their optical and optical--NIR color distributions (as in some giant ellipticals), the apparent inconsistencies most likely result from nonlinearities in the color--metallicity relations.} ",
        "introduction": "The study of globular cluster (GC) systems in galaxies is one of the keystones for understanding the processes at the base of the formation and evolution of galaxies \\citep{ashman92,brodie06}. Recently, the interpretation of one of the most intriguing properties of GC systems in early-type galaxies, the nearly universal presence of two distinct peaks in the optical color distribution, has inspired a vigorous and prolific debate \\citep[][]{yoon06,kundu07,yoon11b,chiessantos12,blake12a,usher12}.  The importance of GC bimodality was recognized before it was a commonly observed property in early-type galaxies \\citep[ETGs,][]{schweizer87}. Historically, the bimodal GC color in optical bands has been equated to metallicity ([Fe/H]) bimodality, implying a fundamental constraint on GC and galaxy formation scenarios. Metallicity bimodality requires two distinct epochs or mechanisms of formation, or both, for the blue (metal-poor) and red (metal-rich) GC subpopulations.  There are various proposed explanations for the GC color bimodality in ETGs: {\\it dissipational} merging of spirals, in which a merger-formed population of red, metal-rich GCs is assumed to appear distinct from the blue, metal-poor GCs of the progenitor spirals \\citep[][]{ashman92}; the {\\it dry} hierarchical assembly, which begins with a massive ``seed'' ETG that has a unimodal metal-rich GC distribution, and in which it is possible to produce a bimodal metallicity distribution through dissipationless accretion of many early-type dwarfs \\citep[][]{cote98}; and the {\\it insitu} formation scenario \\citep{forbes97}. Most of these proposed mechanisms, though, have assumed a simple linear conversion between [Fe/H] and color, which seemed justified from the small fractional age variations among the GCs \\citep[][]{cohen98,kuntschner02,puzia05}.  However, this assumption became the subject of debate when three independent works, using observations and stellar population models, pointed out non-negligible nonlinearities in the color-metallicity relations of GCs \\citep{peng06,richtler06,yoon06}. In particular, Yoon and colleagues and \\citeauthor{richtler06} demonstrated that these nonlinearities naturally produce bimodal color histograms from nonbimodal [Fe/H] distributions. This interpretation, dubbed the projection effect, provided an alternative explanation based on stellar evolution for the ubiquity of bimodal GC color distributions.  In this regard, \\citet{cantiello07d} suggested the use of multicolor GC histograms to verify the consistency of [Fe/H] distributions derived from different colors. These authors highlighted the role of optical to near-infrared (NIR) colors to distinguish between genuine bimodality in [Fe/H] and projected bimodality in color.  If the nonlinear projection is at work, then the [Fe/H] distributions inferred from linear inversion of different color indices for the same GC sample will show some degree of inconsistency (discordant [Fe/H] peaks and/or fractions of GCs in each [Fe/H] component). The analysis of optical-to-NIR GC colors in various ETGs, as well as $u$ to $z$ photometry in some Virgo cluster members, indicates that the nonlinear projection effect is present at some level in these galaxies \\citep{blake12a,chiessantos12,yoon11b,yoon13}.  At the same time, the results of the SLUGGS survey\\footnote{\\url{http://sluggs.ucolick.org.}}, which collected spectra of $\\sim1000$ GCs in 11 galaxies and derived [Fe/H] from the calcium II triplet index, CaT, support true [Fe/H] bimodality in at least some galaxies in addition to the \\object{Milky Way}.  \\citet{usher12} found evidence for bimodal CaT distributions in six of eight galaxies with sufficient numbers of GC spectra.  Nevertheless, the spectroscopically and photometrically derived [Fe/H] distributions show non-negligible differences in several galaxies of the SLUGGS sample, thus lacking the aforementioned multi-index coherence. However, the case of \\object{\\object{NGC\\,3115}}, an isolated lenticular galaxy at a distance of $\\sim10$~Mpc \\citep[][]{tonry01}, revealed highly consistent [Fe/H] and color distributions, leading the authors to present this galaxy as a critical test of [Fe/H] bimodality \\citep{brodie12}.  Previous optical VLT FORS2 spectroscopy for 17 GCs in \\object{NGC\\,3115} showed hints of both a bimodal metallicity distribution \\textit{and} color or spectral index nonlinearity \\citep[][]{kuntschner02}, but it was limited by sample size.  In this Letter, we combine new NIR photometry with literature data to investigate the consistency of the bimodality in the optical, NIR, optical-NIR colors, and CaT, of the GCs in \\object{NGC\\,3115}. ",
        "conclusions": "Our analysis of the GC system in \\object{NGC\\,3115} provides additional definitive proof of GC [Fe/H] bimodality in this S0 galaxy, the first to be firmly established beyond the \\object{Milky Way}. Various studies have noted that to rule out the possibility that nonlinearities project a nonbimodal [Fe/H] distribution into a bimodal color distribution, one must recover consistent [Fe/H] distributions from multiple different photometric indices.  In particular, \\citet{cantiello07d} showed that optical-to-NIR colors are the most useful in constraining the underlying metallicities.  This is in part because the broad color baselines imply a lower sensitivity to the detailed shape of the index-metallicity relation; for the converse reason, bimodal [Fe/H] distributions may not be evident in the purely NIR colors.  The present study confirms these expectations: $i)$ the optical and optical-to-NIR colors are clearly bimodal and consistent with each other in terms of proportions of red and blue GCs, and $ii)$ the pure NIR distributions give ambiguous results, in the sense that different colors and/or statistical indicators indicate the presence or lack of bimodality.  The comparison of GC colors with SSP models confirms earlier results \\citep{brodie12}, using CaT as a proxy for metallicity, which derived a bimodal [Fe/H] distribution with a dip at $-1.3\\lsim [Fe/H] \\lsim -0.3$ dex.  The model comparison also suggests that the metal-rich GCs have slightly younger ages than the blue/metal-poor ones. Although this conclusion is subject to model uncertainties, the size of the age difference is consistent with that found between the metal-rich and metal-poor GCs in the \\object{Milky Way} \\citep{vandenberg13}.  Some previous studies of optical-to-NIR GC colors in giant ellipticals have found significantly different optical versus optical-to-NIR color distributions.  For instance, using HST ACS and WFC3/IR data, \\citet{blake12a} showed that while the optical color distributions of the Fornax giant elliptical \\object{NGC\\,1399} are clearly bimodal \\citep[supporting earlier results from][]{forte05}, the $V{-}H$ and $I{-}H$ distributions are not, or they imply significantly \\textit{different} bimodal breakdowns than found for the optical alone.  Similarly, \\citet{chiessantos12} studied the $gzK$ color distributions of 14 elliptical galaxies, mainly in the Virgo cluster, and reported that double-peaked color distributions are more common in $g{-}z$ than in the optical-NIR colors.  Both studies found significant nonlinearity between the purely optical and optical-to-NIR colors, and both concluded that bimodal optical color distributions are not necessarily indicative of underlying bimodality in metallicity.  The lack of consistency between the purely optical and optical--NIR colors for \\object{NGC\\,1399} and some Virgo ellipticals contrasts strongly with the coherent color bimodality observed in \\object{NGC\\,3115}, for which the metallicities are also clearly bimodal.  In light of this contrast, and with the evidence for nonlinear color--color relations, the inconsistent optical and optical-NIR color bimodalities in some galaxies imply that nonlinearities do indeed play an important role in shaping the GC color distributions in those galaxies.  Thus, while our multicolor photometric analysis confirms the [Fe/H] bimodality of \\object{NGC\\,3115} GCs, we conclude that the nonlinear projection effect remains a viable explanation for the \\textit{ubiquity} of optical color bimodality and is the most likely cause in cases where the optical-NIR colors lack obvious bimodality.  In summary: $i)$ optical and optical-to-NIR colors and CaT indices of GCs in \\object{NGC\\,3115} are bimodal; $ii)$ the bimodal distributions derived for different photometric and spectroscopic indices show good consistency; $iii)$ evidence for bimodality is weak or absent for purely NIR colors; $iv)$ our results agree with model predictions for GC systems with truly bimodal [Fe/H] distributions, which provides definitive proof of [Fe/H] bimodality in \\object{NGC\\,3115}, perhaps the first galaxy beyond the Local Group for which this is the case; $v)$ comparison with SSP models confirms earlier metallicity results based on CaT indices; $vi)$ despite the consistency of the color distributions, we also observe color-color nonlinearities, most clearly in the case of $i{-}H$ versus $g{-}z$.  Thus, the metallicity distributions of extragalactic GC systems and, more specifically, the existence or lack of a universal bimodality of [Fe/H] in ETGs, remains a matter of debate.  Indirectly, when taken in the broader context of previous work, our analysis indicates that optical GC color bimodalities have different causes in different galaxies, with nonlinear color-metallicity relations playing an important role for some previously studied giant ellipticals. Finally, our study shows the effectiveness of optical-to-NIR colors as an unambiguous test for underlying metallicity bimodality."
    },
    "1402/1402.6329_arXiv.txt": {
        "abstract": "We present observations of the Auriga-California Molecular Cloud (AMC) at 3.6, 4.5, 5.8, 8.0, 24, 70 and 160 \\micron\\ observed with the IRAC and MIPS detectors as part of the {\\it Spitzer} Gould Belt Legacy Survey. The total mapped areas are 2.5 deg$^2$ with IRAC and 10.47 deg$^2$ with MIPS. This giant molecular cloud is one of two in the nearby Gould Belt of star-forming regions, the other being the Orion A Molecular Cloud (OMC). We compare source counts, colors and magnitudes in our observed region to a subset of the SWIRE data that was processed through our pipeline. Using color-magnitude and color-color diagrams, we find evidence for a substantial population of 166 young stellar objects (YSOs) in the cloud, many of which were previously unknown. Most of this population is concentrated around the \\lkha 101 cluster and the filament extending from it. We present a quantitative description of the degree of clustering and discuss the fraction of YSOs in the region with disks relative to an estimate of the diskless YSO population.  Although the AMC is similar in mass, size and distance to the OMC, it is forming about 15 -- 20 times fewer stars. ",
        "introduction": "\\label{sec:intro}  The cycle 4 {\\it Spitzer Space Telescope} Legacy project ``The Gould Belt: Star Formation in the Solar Neighborhood'' (PID: 30574; PI: L.E. Allen) completed the {\\it Spitzer} survey of the large, nearby star-forming regions begun by the c2d Legacy Project \\citep{Evans2003, Evans2009}.  The cloud with the least prior study included in the survey is the cloud we have designated as ``Auriga'' which lies on the Perseus-Auriga border. This cloud has also been designated the California Molecular Cloud by \\cite{Ladaetal2009} since it extends from the California Nebula in the west to the \\lkha 101 region and associated NGC 1529 cloud in the east. We adopt the name Auriga-California Molecular Cloud (AMC) to encompass both nomenclatures.  Despite the AMC's proximity to two of the most well-examined star-forming clouds, Taurus-Auriga and Perseus, it is a relatively unstudied region. Several dark nebulae were noted along its length by \\cite{Lynds1962}, and CO associated with many Lynds objects was measured by \\cite{Ungerechts1987}, who note the presence of a CO ``cloud extending from the California nebula (NGC 1499) in Perseus along NGC 1579 and \\lkha 101 well into Auriga'' (their cloud 12). Only very recently has a giant molecular cloud been unambiguously associated with the series of Lynds nebulae through high resolution extinction maps by \\cite{Ladaetal2009} who placed its distance firmly within the Gould Belt (GB) at $450 \\pm 23$ pc.  At this distance, the cloud's extent of 80 pc and mass of $\\sim10^5 \\ M_{\\odot}$ rivals that of the Orion Molecular Cloud (L1641) for the most massive in the Gould Belt.  For the remainder of this paper, we adopt this distance of 450 pc for the entire AMC. This is consistent with the distance of $510^{+100}_{-40}$ pc found by \\citep{Wolketal2010} on their study of \\lkha 101 with Chandra. We note that this distance differs from that adopted by \\cite{Gutermuthetal2009} for \\lkha 101  of 700 pc.  We have mapped a significant fraction of the AMC with the Infrared Array Camera (IRAC; \\citealt{Fazio2004}) and the Mid-Infrared Photometer for \\Spitzer\\ (MIPS; \\citealt{Rieke2004}) on board the \\textit{Spitzer Space Telescope} \\citep{Werner2004}, with a total overlapping coverage of 2.5 deg$^2$ in the four IRAC bands (3.6, 4.5, 5.8 and 8.0 \\micron) and 10.47 deg$^2$ in the three MIPS bands (24, 70, and 160 \\micron). The mapped areas are not all contiguous and were chosen to include the areas with \\Av~$> 3$, as given by the \\cite{Dobashi2005} extinction maps. The goal of these observations is to identify and characterize the young stellar object (YSO) and substellar object populations. The data presented here are the first mid-IR census of the YSO population in this region. The area around \\lkha 101 and its associated cluster was observed as part of a survey of 36 clusters within 1 kpc of the Sun with \\Spitzer\\ by \\cite{Gutermuthetal2009} and those data have been incorporated into our dataset through the c2d pipeline.  More recently, the AMC has been observed by the {\\it Herschel Space Observatory} at 70 -- 500 \\micron, and by the Caltech Submillimeter Observatory with the Bolocam 1.1 mm camera \\citep{Harveyetal2013}. These observations characterize the diffuse dust emission and the cooler Class 0 and Class I objects which can be bright in the far-IR. We do not analyze the large scale structure of the cloud in this paper as \\cite{Harveyetal2013} present such an analysis with the \\Herschel~observations, which are more contiguous and have a higher resolution than our MIPS observations. \\cite{Harveyetal2013} also include a comparison to these MIPS data and so further analysis is not required here.  We describe the observations and data reduction (briefly as it is well-documented elsewhere) in $\\S$ \\ref{sec:obs}. In $\\S$ \\ref{sec:yso}, we describe the source statistics, the criteria for identifying and classifying YSO candidates and we compare the YSO population to other clouds. The SEDs and disk properties of YSOs are modeled in $\\S$ \\ref{sec:sed}. We characterize the spatial distribution of YSOs in $\\S$ \\ref{sec:spatial} and summarize our findings in $\\S$ \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We observed the AMC with IRAC and MIPS aboard the \\textit{Spitzer Space Telescope} and identify 138 YSOs in the cloud. As our IRAC coverage is segmented, we complemented our more contiguous MIPS coverage with WISE data to further eliminate galaxies from the sample, leaving 28 MIPS-only YSOs remaining, bringing the total number of YSOs in the AMC to 166. We classified the YSOs based on the spectral slope of their SEDs between 2 \\micron\\ and 24 \\micron\\ and find 37 Class I objects, 21 Class F objects (flat spectrum sources), 91 Class II objects, and 17 Class III objects. The high fraction of Class Is and Class Fs suggests that the AMC is relatively unevolved compared to other star-forming clouds. Despite the similarity in cloud properties between the AMC and the OMC, there is a distinct difference in the star formation properties. The star formation in the AMC is also concentrated along its filament, however, it is also forming a factor of about 20 fewer stars than the OMC.  \\cite{Ladaetal2009} find that there is much less material at high density in the AMC than in the OMC and attribute the difference in star formation to this. Further studies of the star formation and YSO population in the AMC are needed to highlight the differences of the two clouds given their similar age.  We modelled the SEDs of the Class II and Class III sources and their excesses by first fitting a K7 stellar spectrum to the optical and near-IR fluxes. The spectrum is normalized to the 2MASS flux (or the IRAC1 flux when 2MASS is unavailable) and we use an \\Av~value to match the spectrum of the stellar model to the de-reddened observed optical fluxes. An A0 stellar spectrum is used in the eight cases where a K7 spectrum is unable to provide a reasonable fit. Fitting a stellar spectrum allows us to measure the disk luminosities and characterize the excess. The excesses of the Class II and Class III sources were further parameterized by \\lambdaturnoff, the longest wavelength before an excess greater that 80\\% is measured, and \\alphaexcess, the slope of the SED at wavelengths longward of \\lambdaturnoff. \\lambdaturnoff\\ is a useful tracer for the proximity of dust to the star and consequently we identify fourteen classical transition disk candidates.  The bulk of the star formation in the AMC is in the southern region of the cloud. We included a clustering analysis to quantify the densest areas of star formation and to identify groups within the cloud. We find four groups with 10 or more members all in the region around \\lkha 101 and its adjoining filament. We find three smaller groups with 5 -- 9 members scattered throughout the cloud. The largest group is that around \\lkha 101 and contains 49 members. We note that there are likely even more YSOs in this group since our YSO identification criteria of S/N $\\geq 3$ in IRAC1-4 and MIPS1 are difficult to attain in this bright region."
    },
    "1402/1402.1215_arXiv.txt": {
        "abstract": "\\begin{changeb} The design of an interstellar microwave digital communication system at interstellar distances is considered, with relevance to starships and extraterrestrial civilizations (SETI and METI). Distances are large and thus require large transmitted power and/or large antennas, while the atmospheric/interstellar microwave is transparent over a wide bandwidth. Recognizing the  known tradeoff between wider signal bandwidth and lower energy, a reduced requirement for energy delivered to the receiver at the expense of wide bandwidth is advantageous. In addition, wide bandwidth results in significant simplifications to the design and implementation. It is shown that appropriate transmit signal design can circumvent dispersion and scattering  arising in the interstellar medium and motion, obviating any related processing in transmitter or receiver. The concept of interstellar coherence hole is introduced, capturing the time and frequency coherence constraint. Because other impairments can be circumvented, the fundamental limit on the energy delivered to the receiver for each bit of information is determined by cosmic microwave background alone. Designs that exchange larger bandwidth for extreme simplicity (recognizing the lack of coordination) and approaching the fundamental energy limit are identified. The Morse code invented for the telegraph in 1836 comes closer to this ideal than most approaches used in modern terrestrial radio by mapping a single bit onto a carrier burst. Rather than the terrestrial approach of adding phases and amplitudes to increase information capacity while minimizing bandwidth, adding multiple locations for carrier bursts in time and frequency increases information capacity while minimizing energy per information bit. The resulting location code can approach the fundamental energy limit as bandwidth is expanded, and is consistent with blind discovery with straightforward modifications to existing SETI pattern recognition algorithms. A discovery search can detect individual carrier bursts absent detailed knowledge of the signal structure, and observations of the interstellar medium and motion by transmitter and receiver desirably constrain the parameterization of these bursts. \\end{changeb} ",
        "introduction": "There has been considerable experience in communication with ''deep space'' probes in and around our solar system at radio wavelengths \\citeref{660}, and there is growing interest in optical \\citeref{661} as well. Communication of information at much greater distances, such as with starships and extraterrestrial civilizations, introduces new challenges, and has not been addressed either theoretically or empirically. Some of those challenges are addressed here from the perspective of communication engineering, emphasizing radio (rather than optical) wavelengths. \\begin{changea} The insights here are highly relevant to both the search for extraterrestrial intelligence (SETI) by informing the discovery of information-bearing signals, and messaging for extraterrestrial intelligence (METI), which transmits information-bearing signals. \\end{changea} \\begin{changec} They are also relevant to the design of two-way links with interstellar spacecraft (often called starships). \\end{changec}  \\doctable {acronyms} {Acronyms} {l p{6cm}} { \\bfseries Acronym & \\bfseries Definition \\\\ \\hline AWGN & Additive white Gaussian noise \\\\ CMB & Cosmic background radiation \\\\ CSP & Coding and signal processing \\\\ ICH & Interstellar coherence hole \\\\ ISM & Interstellar medium \\\\ SETI & Search for interstellar intelligence \\\\ \\hline }  \\begin{changea} This context is illustrated in Fig. \\ref{fig:blockDiagram}, showing the subsystems of an end-to-end system communicating information by radio. This paper concerns the coding and signal processing (CSP) subsystem. It is presumed that a message input is represented digitally (composed of discrete symbols), in which case it can always be represented as a stream of bits. The transmit-end CSP inputs information bits and outputs a baseband waveform \\begin{changec} (the spectrum of which is concentrated about d.c.). \\end{changec} The radio subsystem \\begin{changec} (including modulation to passband, centered about a carrier frequency $f_c$, \\end{changec} a radio-frequency transmitter, transmit/receive antennas, and demodulation \\begin{changec} back to baseband) \\end{changec} delivers a replica of this baseband waveform to the receive-end CSP with impairments (such as distortion and noise) introduced by the physical environment (the interstellar propagation and motion). The modeling of these impairments, which profoundly affect the CSP, is summarized here and addressed in greater depth elsewhere \\citeref{664}.  \\incfig {blockDiagram} {1} {The major functional blocks in an interstellar communication system, with transmit on the top row and receive on the bottom row. These blocks are separated into (from left to right) information, baseband processing, passband radio, and physical propagation. }  The CSP realizes the essential function of mapping information bits into a continuous-time baseband waveform suitable for transmission as an electromagnetic wave. The radio subsystem attempts to minimize impairments, for example by reducing noise introduced in the receiver. Through the choice of antennas, in the context of a given transmission distance it determines the radiated power necessary to deliver a needed level of power at receive baseband. The remaining performance characteristics of the system are determined by the CSP, including the fidelity of the message replica delivered by the receive-end CSP and the resources consumed in achieving this fidelity. The fidelity is usually measured by probability of error, and the primary resources of interest are the bandwidth of the radiated signal and the signal power that must be delivered to receive baseband. These resources can be manipulated over many orders of magnitude through the design of the CSP, and thus it is the subsystem with the greatest opportunity to manipulate the message fidelity and resources.  A complete messaging system combines the expertise of several disciplines, including astronomy and astrophysics to model the physical, radio astronomy and electronics to design the radio, and applied statistics to design the CSP. The CSP has been the target of extensive research and practice in the context of commercial terrestrial wireless systems, and that experience is directly applicable to the interstellar case. The goal here is to communicate to a wider audience the opportunities afforded by the CSP. Since the knowledge and techniques applied in the CSP may be unfamiliar to many readers, the emphasis here is on reviewing results and developing intuition, with details documented elsewhere \\citeref{656}.  One of the tradeoffs determined by the CSP is between bandwidth and received power; inevitably, increasing either allows the other to be reduced. \\end{changea} We argue that in communication at the large interstellar distances, large propagation losses place a premium on minimizing energy requirements. Other challenges applying specifically to communication with other civilizations are dispersive effects arising in the interstellar medium (ISM) at distances greater than hundreds of light years, relative motion, a lack of prior design coordination between the two ends, and the challenges of discovering a signal with unknown design and parameterization. It turns out that these considerations are closely linked, with both energy minimization and interstellar impairments contributing beneficially to both an implicit form of coordination between transmitter and receiver, as well as informing a discovery search. \\begin{changea} In the case of starships, relativistic effects also become critical due to large relative velocities \\citeref{665}. \\end{changea} ",
        "conclusions": "End-to-end design of a digital communication system at interstellar distances has been considered, with an emphasis on minimizing the energy per bit or energy per message delivered to the receiver. A solution with all the desired properties is identified. \\begin{changeb} A location code is conceptually simple, and a signal of this type is straightforward to discover at the receiver using a strategy based on detection of individual energy bursts followed by pattern recognition. \\end{changeb} This design can also approach the bandwidth-unconstrained limit on received energy, even in the presence of scintillation (when it is combined with time diversity). It does, however, require a significant bandwidth expansion. A similar approach can also achieve energy-efficient information-free beacons \\citeref{656}.  Perhaps most significantly, the fundamental principle of minimizing received energy without bandwidth constraint and in consideration of jointly observable ISM and motion impairments leads to a simple and highly-constrained signal design. One can hope that the transmitter and receiver designers addressing this joint challenge without the benefit of coordination might arrive at compatible conclusions. Interstellar impairments are fortuitous for uncoordinated communication since they constrain the signal structure and parameterization. The tighter constraints at lower carrier frequencies and at radio as opposed to optical wavelengths are also helpful in this regard, because they more tightly constrain the signal without adversely impacting the energy requirements.  It is reasonable to ask two skeptical questions. First, is it likely that another civilization is aware of the opportunity to reduce received energy, and aware of the fundamental limit on received energy? The history of earthbound communications leads to optimism. Energy-limited communication near the fundamental limit was understood much earlier (in the 1950`s) than bandwidth-limited communication (in the 1990`s) because of the simplicity of the solution. Optical communication is typically much simpler than radio precisely because bandwidth has never been considered a limiting resource at the shorter optical wavelengths.  Second, is another civilization likely to be motivated by and act upon the opportunity to reduce the received energy? This is a more difficult question, since a more advanced civilization may well have tapped into cheaper sources of energy. Even if so, they may be motivated by the simplicity of the solutions and the benefits of that simplicity to implicit coordination. In addition, even if energy is more plentiful, there are many beneficial ways to consume more energy other than deliberate inefficiency. They could increase message length, reduce the message transmission time, transmit in more directions simultaneously, or transmit to greater distances. Overall it is unlikely that a civilization would use more energy than necessary unless for some reason they consider a reduction in bandwidth to be a higher priority.  For communication with a starship, trading for greater bandwidth remains useful as a way to minimize received energy, which is a particular benefit on the downlink from the starship because of the reduction in transmit power and/or transmit antenna size. Because there is the luxury of joint design of transmitter and receiver, simplicity is not a goal in itself and if desired the bandwidth expansion can be reduced without a significant penalty in received energy using something considerably more complex than the location code."
    },
    "1402/1402.0814_arXiv.txt": {
        "abstract": "{ The treatment of radiation transport in global circulation models (GCMs) is crucial for correctly describing Earth and exoplanet atmospheric dynamics processes. The two-stream approximation and correlated-$k$ method are currently state-of-the-art approximations applied in both Earth and hot Jupiter GCM radiation schemes to facilitate the rapid calculation of fluxes and heating rates. Their accuracy have been tested extensively for Earth-like conditions, but verification of the methods' applicability to hot Jupiter-like conditions is lacking in the literature. We are adapting the UK Met Office GCM, the Unified Model (UM), for the study of hot Jupiters, and present in this work the adaptation of the Edwards--Slingo radiation scheme based on the two-stream approximation and the correlated-$k$ method. We discuss the calculation of absorption coefficients from high-temperature line lists and highlight the large uncertainty in the pressure-broadened line widths. We compare fluxes and heating rates obtained with our adapted scheme to more accurate discrete ordinate (DO) line-by-line (LbL) calculations ignoring scattering effects. We find that, in most cases, errors stay below $\\SI{10}{\\percent}$ for both heating rates and fluxes using $\\sim 10$ $k$-coefficients in each band and a diffusivity factor $D = 1.66$. The two-stream approximation and the correlated-$k$ method both contribute non-negligibly to the total error. We also find that using band-averaged absorption coefficients, which have previously been used in radiative-hydrodynamical simulations of a hot Jupiter, may yield errors of $\\sim \\SI{100}{\\percent}$, and should thus be used with caution.} ",
        "introduction": "For Earth's atmosphere, irradiation from the Sun is the primary source of energy. Any model of the Earth's atmosphere therefore needs a robust and accurate treatment of radiation transport. Global circulation models (GCMs) of the Earth are used for both weather prediction and climate research, and they include a dynamical core that solves some variant of the Navier--Stokes equations and a radiation scheme that calculates the radiative heating rate. The dynamical cores are tested using benchmarks \\cite[see e.g.][]{Held1994,Reed2011,Ullrich2013}, and both dynamical cores\\footnote{See \\url{http://earthsystemcog.org/projects/dcmip-2012/}} and radiation schemes \\citep{Ellingson1991,Collins2006,Oreopoulos2012} are tested through intercomparison projects.  GCMs have also been successfully applied to other solar system planets such as Jupiter, Saturn, Mars and Venus \\citep[see for example][respectively]{Yamazaki2004, Muller-Wodarg2006, Hollingsworth2010, Lebonnois2011}. In the past decade, GCMs have been used to study large-scale circulation on hot Jupiters \\citep{Showman2002,Showman2009,Rauscher2012,Cho2008,Thrastarson2010,Dobbs-Dixon2008}, a class of extrasolar planets which are approximately the size of Jupiter but orbit less than $\\SI{0.1}{\\astronomicalunit}$ from their parent star. These planets, thought to have tidally locked circular orbits due to strong tidal interactions between the planet and its parent star \\citep{Baraffe2010}, experience intense irradiation yielding a significant temperature contrast between the (permanent) day-side and night-side. Winds in the atmosphere of these planets are therefore expected to transport heat from the day-side to the night-side.  Atmospheric properties of hot Jupiters are obtained by various observational techniques, mainly transmission spectroscopy and secondary eclipse measurements. Brightness maps \\citep{Knutson2007,Majeau2012} and wind velocities \\citep{Snellen2010} are now accessible, and constraints on the composition are becoming available, but large uncertainties remain. Observations indicate a hotspot shifted eastward of the substellar point \\citep{Knutson2007,Majeau2012} and temperature contrasts smaller than what is expected for these planets without winds \\citep{Knutson2007,Knutson2009}, indicating transport of heat from the day-side to the night-side \\citep{Watkins2010,Perez-Becker2013}. HD~209458b appears to have a temperature inversion in its upper atmosphere \\citep{Knutson2008,Burrows2007} while HD~189733b does not \\citep{Charbonneau2008,Barman2008,Knutson2009}, indicating that, despite similar orbital properties, hot Jupiters may have very different circulation patterns that still need to be understood. GCMs are therefore very valuable when trying to understand the increasing amount of observations of these systems.  Benchmarking of the dynamical cores of GCMs applied to hot Jupiters has been used to investigate stability of the codes and discrepancies between them \\citep{Heng2011,Menou2009,Bending2013,Mayne2013a,Mayne2013b,Polichtchouk2014}. These benchmarks, and early GCMs applied to hot Jupiters, used simple, parametrised radiation schemes termed ``Newtonian cooling'' or ``temperature forcing'', where the temperature is relaxed towards assumed equilibrium pressure--temperature ($P$--$T$) profiles on a given timescale \\citep{Showman2002}. Pressure--temperature profiles and timescales can be estimated using one-dimensional time-dependent radiative transfer calculations \\citep{Iro2005}, though such an approach has flaws: (i) The equilibrium $P$--$T$ profiles used in the forcing may have a limited accuracy, (ii) radiative timescales may also have a limited accuracy and will vary in a non-trivial way as a function of latitude, longitude and depth, (iii) the forcing parametrisation itself may not be physically realistic, though the use of time-averaged equilibrium states when analysing model results may make this less of an issue, (iv) the model flexibility is poor since for each new planet modelled, the forcing must be changed.  Later studies used more complicated schemes such as flux-limited diffusion \\citep{Dobbs-Dixon2008} and the two-stream approximation \\citep{Showman2009,Rauscher2012,Dobbs-Dixon2013}. For the opacity treatment, grey schemes \\citep{Rauscher2012}, binning and averaging of the absorption coefficients \\citep{Dobbs-Dixon2013} and the correlated-$k$ method \\citep{Showman2009} have been used. The correlated-$k$ method has also been used for retrieval analysis and characterisation of hot Jupiter atmospheres \\citep{Irwin2008} and to model brown dwarf atmospheres \\citep{Burrows1997}. Brown dwarfs atmospheres have many similarities with hot Jupiter atmospheres (e.g. temperature range and composition), but local conditions are very different due to the strong irradiation from the parent stars for hot Jupiters. There is a notable lack of analysis of the accuracy of these schemes when applied to hot Jupiter-like atmospheres and of details on how opacities have been calculated from line lists, preventing rigorous comparison with results previously published in the literature. These are serious shortcomings in a field of research which develops quickly and will deliver more and more accurate data requiring reliable tools for their interpretation.  Both the two-stream approximation and the correlated-$k$ method \\citep[see e.g.][]{Thomas2002} are widely used in GCM simulations of the Earth, and the literature on the methods' applicability to the Earth atmosphere and their accuracy is extensive \\citep[see e.g.][]{Toon1989,Meador1980,Zdunkowski1980,Goody1989,Lacis1991,Mlawer1997}. They have both been found to yield results with satisfactory accuracy when comparing to more accurate solutions obtained from e.g. discrete ordinate (DO), line-by-line (LbL) calculations and when different schemes are compared through intercomparison projects \\citep{Ellingson1991,Collins2006,Oreopoulos2012}. They are, however, still under investigation \\citep{Goldblatt2009} and are still one of the limiting factors of the accuracy of both weather prediction and climate modelling.  We are currently adapting the UK Met Office GCM, the Unified Model (UM), to hot Jupiter-like conditions. The main advantage of this model is its dynamical core, which solves the full 3D Euler equations and is coupled to a radiation scheme based on the correlated-$k$ method and two-stream approximation. This GCM is state-of-the-art for both Earth and hot Jupiter atmospheric dynamics modelling. In previous papers we have tested and confirmed the dynamical core's suitability to model hot Jupiter-like atmospheres \\citep{Mayne2013a,Mayne2013b}. This paper presents the adaptation of the radiation scheme, the Edwards--Slingo (ES) radiation scheme \\citep{Edwards1996a}, to conditions prevailing in hot Jupiter atmospheres. Opacities from high-temperature line lists have been calculated for the dominant absorbers in these atmospheres and the radiation scheme, using $k$-coefficients calculated from these opacities, has been tested against more accurate DO LbL calculations.  Observations of absorbing and scattering species in hot Jupiter atmospheres have so far been limited to the detection of molecular absorbers \\citep{Huitson2013,Wakeford2013,Tinetti2007}, with some observations suggesting Rayleigh/Mie scattering clouds \\citep{Pont2008,Sing2013}. Due to the large uncertainties related to scatterers in hot Jupiter atmospheres and the complexity it adds to radiation transport, we limit the discussions in this paper to purely absorbing atmospheres and postpone the inclusion of scattering to a future work.  The motivation for the present work is the lack of accurate tests and analysis of these radiation schemes now widely used by the community. This work will help in the implementation of similar schemes in the future and provide some guidelines for further progress. In \\cref{sec:opacities} we first discuss our opacity data including the high-temperature line lists we use and the calculation of cross-sections from these line lists by using estimates for the line widths. In \\cref{sec:ES} we briefly summarise the implementation of the correlated-$k$ method and two-stream approximation in the UM and move on to testing this scheme for gradually more complicated hot Jupiter-like atmospheres in \\cref{sec:testing}. These tests will be useful to other groups for comparison and benchmark purposes. Our conclusions follow in \\cref{sec:conclusions}. Combination of the adapted dynamical core and radiation schemes is in progress and will be presented in a future work, which will result in a state-of-the-art GCM that can be applied to a variety of exoplanet atmospheres. ",
        "conclusions": "\\label{sec:conclusions}  The accuracy of radiation schemes used in GCMs has been studied extensively for Earth-like conditions, but detailed analysis for hot Jupiter-like conditions are lacking. In this paper we have analysed the accuracy and uncertainties in state-of-the-art radiation schemes used in several GCMs applied to hot Jupiters. Opacity sources and calculation of absorption coefficients from high-temperature line lists have been discussed. We present a line profile cut-off scheme that decreases the computation time required to calculate absorption coefficients by a factor of $\\sim 100$ compared to other methods used in the literature, while still giving accurate results. Both the two-stream approximation and correlated-$k$ method's applicability to hot Jupiter atmospheres have been analysed by comparing the Edwards--Slingo radiation scheme to discrete ordinate line-by-line calculations.  The ES radiation scheme's performance in these tests shows that we have successfully adapted it to hot Jupiter-like atmospheres. Our main conclusions are: \\begin{itemize} \\item Pressure broadening parameters for high-temperature molecular lines are very uncertain and usually extrapolated from room temperature and pressure and small quantum numbers. Improvements in this area will become important as higher accuracy will be required to analyse results from future exoplanet characterisation projects (e.g. JWST, EChO, Sphere and ELT). \\item A diffusivity factor of $D = 1.66$, already widely used in both Earth and hot Jupiter GCMs, yields the smallest errors from the two-stream approximation, although $D = \\sqrt{3} \\approx 1.73$ is only slightly less accurate. \\item About $10$ $k$-coefficients in each band for molecular line absorption yield satisfactory accuracy. Using $\\sim 100$ $k$-coefficients per band does improve the overall accuracy, but errors decrease by less than $\\SI{50}{\\percent}$, while the radiative transfer computation time increases by a factor of $10$. We therefore choose to adopt the former as a balance between accuracy and computational cost. \\item Both the two-stream approximation and the correlated-$k$ method contribute non-negligibly to the total error, with overall heating rate errors of $\\lesssim \\SI{10}{\\percent}$ in regions with significant heating/cooling. Flux errors are similar or smaller. \\item Whether a black-body spectrum, solar spectrum or uniform (in wavenumber) weighting scheme is used has little effect on the overall accuracy given the band structure used here (\\cref{tbl:bands}). We therefore choose to adopt a uniform weighting scheme, enabling the use of the same $k$-coefficients in both the thermal and stellar spectral regions and for different irradiation spectra. \\item Using a mean absorption coefficient in each band, as in \\citet{Dobbs-Dixon2013}, yields inaccurate fluxes and heating rates for molecular absorption. Heating rate errors can reach $\\SI{100}{\\percent}$ or more, even in regions with significant heating. Band-averaged absorption coefficients should thus be used with caution. \\end{itemize}  Any radiation scheme applied to hot Jupiters should be checked against the tests we have presented here. These tests and the detailed descriptions of our methods and approximations will be useful for future adaptation of radiation schemes in other GCMs. Current observational constraints on exoplanets do not require the level of accuracy we have applied in this work. The field develops at an amazing pace, however, and modellers should now develop the best theoretical and numerical tools to tackle the challenges posed by the increasing accuracy expected from future large observational projects."
    },
    "1402/1402.2485_arXiv.txt": {
        "abstract": "In order to determine the composition of the dust in the circumstellar envelopes of oxygen-rich asymptotic giant branch (AGB) stars we have computed a grid of {\\sc modust} radiative-transfer models for a range of dust compositions, mass-loss rates, dust shell inner radii and stellar parameters. We compare the resulting colours with the observed oxygen-rich AGB stars from the SAGE-Spec Large Magellanic Cloud (LMC) sample, finding good overall agreement for stars with a mid-infrared excess. We use these models to fit a sample of 37 O-rich AGB stars in the LMC with optically thin circumstellar envelopes, for which 5--35-\\mum {\\em Spitzer} infrared spectrograph (IRS) spectra and broadband photometry from the optical to the mid-infrared are available. From the modelling, we find mass-loss rates in the range $\\sim 8\\times10^{-8}$ to $5\\times10^{-6}$ M$_{\\odot}\\ \\mathrm{yr}^{-1}$, and we show that a grain mixture consisting primarily of amorphous silicates, with contributions from amorphous alumina and metallic iron provides a good fit to the observed spectra. Furthermore, we show from dust models that the {\\em AKARI} [11]--[15] versus [3.2]--[7] colour-colour diagram, is able to determine the fractional abundance of alumina in O-rich AGB stars. ",
        "introduction": "Stars on the asymptotic giant branch (AGB) lose a significant fraction of their matter through slow dense winds at rates of $10^{-10}$ to $10^{-4}$ M$_\\odot$ yr$^{-1}$ \\citep{Bowen1991}. During this phase of evolution the effective stellar temperature is low enough for molecules to form and dust grains to condense in the circumstellar outflow. In general, dust formation occurs in stages known as the `dust condensation sequence', and depends on the physical conditions in the envelope \\citep[][]{Tielens1990, GailSed1999, Gail2010}. The mineralogy of the dust is dominated by the chemistry of the envelope (which is set by a star's photospheric C/O atomic ratio and initial elemental abundance), and by the envelope's temperature and density (which determines which condensation products are stable).  Although not understood in detail, the thermodynamic condensation sequence for oxygen-rich outflows (C/O $<$ 1) predicts that refractory oxides such as alumina (Al$_{2}$O$_{3}$) would be the first astrophysically significant species to form. These grains can exist relatively close to the star ($T_{\\rm cond} \\sim 1400$ K at pressures relevant to stellar outflows) and may act as seed nuclei for other grains \\citep[e.g.][]{Onaka1989, Stencel1990, Sogawa1999}. At larger radii, where the temperature is lower ($T \\sim 1000$ K) silicate grains form.  It is well established that the dust in oxygen-rich outflows consists mainly of silicates  \\citep{WoolfNey1969, Hackwell1972, TreffersCohen1974}. However, alumina dust has been detected in the spectra of some low mass-loss rate ($\\dot M$) O-rich AGB stars in the Milky Way \\citep{Onaka1989, Speck2000, Dijkstra2005}. These stars show a broad low-contrast emission feature in the 7 -- 14 $\\mu$m region which is best fit with a blend of amorphous alumina and amorphous silicates. As the star evolves along the AGB the shape and peak position of this feature changes; silicate dust becomes more important and dominates the 10-$\\mu$m emission \\citep{Little-Marenin1990,Sloan1995}. At high mass-loss rates the 10-$\\mu$m feature is only due to silicates. This sequence has been quantified by \\cite{Sloan1995, Sloan1998} using flux ratios in the 9--12 $\\mu$m region.  The \\emph{Spitzer} Space Telescope has taken spectra of a large number of individual evolved stars in the Large and Small Magellanic Clouds, covering a wide range of colours and magnitudes \\citep{Kemper2010}. Although the mineralogy of the oxygen-rich AGB stars in the Magellanic Clouds is dominated by amorphous silicates both at low and high mass loss rates \\citep{Sargent2010, Riebel2012} some sources also contain a small fraction of crystalline silicates \\citep{Jones2012}. There is an apparent absence of Al$_{2}$O$_{3}$ in the \\emph{Spitzer} spectra of O-AGB stars in the Magellanic Clouds \\citep{Sloan2008}, with the 10-$\\mu$m feature of low density (low $\\dot M$) winds better reproduced by amorphous silicate grains than the expected mixture of Al$_{2}$O$_{3}$ and silicates. This implies that either alumina dust is depleted in the LMC or that observational biases hinder its detection. However, an extensive systematic study into the Al$_{2}$O$_{3}$ content of evolved stars in the LMC has not yet been carried out.  In this study we evaluate the alumina content of oxygen-rich AGB stars in the Magellanic Clouds, by creating a grid of radiative transfer models that explores a range of physical parameters relevant to evolved stars. In Section~\\ref{model}, we describe how the radiative transfer models for the circumstellar dust shell were computed. Colour-colour diagrams of the grid of models are presented in Section~\\ref{results} and will be used to study the alumina dust in a sample of O-rich AGB stars in the Magellanic Clouds in Section~\\ref{AluminaLMC}. Finally, we discuss our results in Section~\\ref{Discussion} and summarise this study in Section~\\ref{conclusions}. ",
        "conclusions": "We have presented a grid of dust radiative transfer models which explores a range of alumina and silicate dust compositions and stellar and dust shell parameters. The models have been used to simultaneously fit the spectra and broadband SED of  37 oxygen-rich AGB stars in the LMC with optically-thin circumstellar envelopes. The mass-loss rates of our sample range from $\\sim 8\\times10^{-8}$ to $5\\times10^{-6}$ M$_{\\odot}\\ \\mathrm{yr}^{-1}$. We find that a combination of amorphous silicates, amorphous alumina and metallic iron provides a good fit to the spectra. All our sources we found to be silicate-rich, though alumina and iron is often present in significant amounts. This dust composition is consistent with the thermodynamic dust condensation sequence for oxygen-rich AGB stars. Furthermore, we show from dust models that the {\\em AKARI} [11]--[15] versus [3.2]--[7] colour, is able to determine the fractional abundance of alumina in oxygen-rich AGB stars."
    },
    "1402/1402.5093_arXiv.txt": {
        "abstract": "% A quantum electrodynamics (QED) correction surface for the simplest polyatomic and polyelectronic system \\hp{} is computed using an approximate procedure. This surface is used to calculate the shifts % to vibration-rotation energy levels due to QED; such shifts have a magnitude of up to 0.25~\\cm{} for vibrational levels up to 15~000~\\cm{} and are expected to have an accuracy of about 0.02~\\cm. Combining the new \\hp\\ QED correction surface with existing highly accurate Born-Oppenheimer (BO), relativistic and adiabatic components % suggests that deviations of the resulting {\\it ab initio} energy levels from observed ones  % are largely due to non-adiabatic effects.  \\pacs{Valid PACS appear here}% ",
        "introduction": "\\emph{Ab initio} studies of diatomic and triatomic systems containing less than ten electrons are nowadays able to produce rotation-vibrational energy levels with better than spectroscopic accuracy, i.e. with errors of less than 1~\\cm. To improve on this  accuracy one needs to account for several small effects which are routinely neglected, including electronic relativistic and adiabatic corrections, as well as --- most notably for this work --- non-adiabatic effects and corrections due to quantum electrodynamics (QED). General discussions of relativistic and QED effects in molecular physics and quantum chemistry can be found in several recent reviews \\cite{Pyykko2012,Liu2013,Autschbach2012,Liu2012,Kutzelnigg2012,Ilias2010,anatomy} and textbooks \\cite{Reiher.book,Dyall.book}. In this study we follow the convention of calling `relativistic effects' corrections to the non-relativistic Schr{\\\"o}dinger equation of second order in the fine-structure constant $\\alpha$ (i.e., all effects correctly described by the many-electron no-pair Dirac-Coulomb-Breit equation), while so-called radiative corrections due to the quantization of the electromagnetic field and appearing in higher powers of $\\alpha$ are referred to as QED effects.  The hydrogen molecular ion H$_2^+$ is the simplest physical system with a rotational-vibrational spectrum and serves as an important benchmark. Rotational-vibrational energy levels for H$_2^+$ were notably presented by Moss \\cite{Moss1999} with an estimated accuracy of $10^{-4}$~\\cm{} and included non-adiabatic, relativistic as well as leading QED corrections. More recent studies have considerably improved the achievable accuracy and, for selected rotation-vibrational transitions, QED corrections up to $\\alpha^5$ have been computed \\cite{Korobov2006,Korobov2008,Zhong2012} leading to uncertainties of about $2\\times 10^{-6}$~\\cm.  Next in terms of size and complexity is the hydrogen molecule \\htwo, for which an accuracy of  $10^{-4}$  \\cm\\ has recently been achieved {\\it ab initio} \\cite{KomasaJCTC2011, Ubachs,Ubachs1} by careful inclusion of non-adiabatic corrections and of QED corrections to order $\\alpha^4$. Studies of H$_2^+$ and H$_2$ represent the current state-of-the-art for calculations of molecular rotational-vibrational energy levels; for larger systems the achievable accuracy is considerably lower.  In particular, for \\hp\\ the highest accuracy achieved so far is 0.10~\\cm\\ for all known energies up to 17~000~\\cm\\ \\cite{jt512}, which is therefore several orders of magnitude worse than for H$_2^+$ and H$_2$. Higher accuracy energy levels are necessary for proper analysis of \\hp\\ experimental spectra. More specifically, about 30 years ago Carrington and co-workers \\cite{CB82,CK84,cm89} measured very dense near-dissociation spectra of \\hp\\ and its isotopologues with an average line spacing of less than 0.01 \\cm; these spectra, which remain unassigned and substantially uninterpreted \\cite{jt157}, clearly require very high accuracy to be analysed from theoretical calculation.  Another source of motivation is provided  by the recent studies by Wu {\\it et al} \\cite{13WuLiLi.H3+} and Hodges  {\\it et al} \\cite{13HoPeJe.H3+}, who have concentrated on high-precision and high-accuracy frequency measurements on the \\hp\\ $\\nu_2$ fundamental band. Measurements were made by both groups at the sub-MHz ($3 \\times 10^{-5}$ \\cm) level but currently do not agree with each other within the claimed uncertainties.  The assigned \\hp\\ experimental data has recently been the subject of an analysis using the MARVEL procedure \\cite{jt412}, producing a comprehensive set of rotation-vibration energy levels \\cite{H3pMARVEL,H3pMARVEL1} which we use for comparison throughout this study.  Given the present experimental situation it is therefore very desirable to improve the accuracy of theoretical \\hp\\ energy levels beyond the 0.1~\\cm\\ level. The main non-relativistic, clamped nuclei Born-Oppenheimer (BO) potential energy surface (PES) from  Pavanello  {\\it et al}  \\cite{jt512,jt526} and the associated relativistic and adiabatic surfaces, all of which we use in this work, are probably sufficiently well-determined to predict energy levels with an accuracy of about $10^{-2}$ \\cm\\  for low-lying levels up to about 15,000~\\cm. % There are currently two factors limiting the accuracy in \\hp\\ to the 0.1~\\cm{} level, namely a proper treatment of, \\emph{i)}, non-adiabatic and, \\emph{ii)}, QED effects.  Non-adiabatic effects in \\hp\\ and its isotopologues are known to affect line positions by up to 1.0~\\cm{} \\cite{jt236} and therefore must be accounted for accurately. Polyansky and Tennyson (PT) \\cite{jt236} introduced a simple model based on the use of fixed, effective vibrational and rotational masses taken from  Moss's \\cite{Moss1993} studies on H$_2^+$; PT were able to improve the accuracy of calculations from 1~\\cm\\ to 0.1~\\cm. Further improvements require more sophisticated treatments of non-adiabatic effects; a step in this direction has been made by  Diniz {\\it et al} \\cite{jt566}, who obtained non-adiabatic rotational-vibrational energies for the $\\nu_2$ band with an accuracy of      0.01    ~\\cm\\ but did not consider higher vibrational states.  The second factor limiting the final accuracy of \\hp\\ energy levels are QED effects. As discussed above, QED effects have been computed accurately for H$_2^+$ \\cite{Moss1993} and H$_2$ \\cite{KomasaJCTC2011,Ubachs,Ubachs1} and have an effect in the region 0.1---0.2~\\cm\\ on the corresponding rotation-vibration energy levels. In the case of \\hp, QED effects have so far been  entirely neglected but must clearly be taken into account to achieve accuracies better than 0.1~\\cm.  Pyykk{\\\"o} \\emph{et al.} \\cite{jt236} suggested a simple scheme for describing leading QED effects in molecules (see section \\ref{section.QED} for details). This scheme has been already applied to the water molecule \\cite{jt236,jt309} --- for which QED corrections % are of the order of 1~\\cm --- and was instrumental in recent studies achieving an accuracy of 0.1~\\cm\\ for levels up to 15~000~\\cm\\ \\cite{jt550} and of 1~\\cm\\ for the dissociation energy \\cite{jt549}. In this study we use the model of Pyykk{\\\"o} \\emph{et al.} \\cite{jt236} to provide a QED correction surface for \\hp. This correction energy surface, when combined with the existing non-relativistic, relativistic and adiabatic surfaces from previous studies \\cite{jt512,jt526} and with a future, accurate treatment of non-adiabatic effects is expected to provide rotation-vibration energy levels with a typical accuracy of 0.01 \\cm.  The paper is organised as follows. Section \\ref{sec.basis-set} presents a comparison of the Born-Oppenheimer PES computed using explicitly correlated Gaussians \\cite{jt512,jt526} and surfaces computed using standard quantum chemistry methods based on full configuration interaction (FCI) and Gaussian basis sets.  We  show that available basis sets provide an accuracy between 0.1 \\cm\\ and 1 \\cm\\ for rotation-vibration energy levels.  Section III compares results of accurate QED calculations for \\htwo\\ \\cite{KomasaJCTC2011,Ubachs,Ubachs1} with our calculations using the approximate method of Pyykk{\\\"o} et al. \\cite{jt236}. QED corrections for \\hp\\ using the  same methodology are presented. Section  IV presents results of nuclear motion calculations using a BO PES, relativistic and adiabatic corrections \\cite{jt512,jt526} and our QED correction surface. Nuclear motion calculations are given both without non-adiabatic corrections and with a simple non-adiabatic treatments based either on the Polyansky-Tennyson (PT) model  \\cite{jt236} or on the model by Diniz {\\it et al} \\cite{jt566}.  Analysis of the residual deviations between theory and experiment is given. Section V presents a final discussion and conclusions. ",
        "conclusions": "We  calculated a QED energy correction surface for \\hp\\ using the approximate method of Pyykk{\\\"o} \\etal \\cite{jt265}.  This method is benchmarked against accurate QED calculations for H$_2^+$ and \\htwo; the comparisons suggest that our QED surface for \\hp\\ should provide QED corrections to rotational-vibration energy levels with an accuracy better than $0.02$~\\cm. The effect of QED on low-lying energy levels is of the order of 0.2~\\cm{} and hence is much larger than the accuracy of $10^{-2}$ \\cm\\ which has already been achieved for all components of \\ai\\ calculations on \\hp\\ with the notable exception of non-adiabatic effects.  Inclusion of QED effects leads to \\hp\\ energy levels being reproduced with a RMS % deviation which is reduced from 0.99 \\cm\\  to 0.84 \\cm{} when no allowance is made for non-adiabatic effects (nuclear masses used for energy levels calculation). These calculations, which include highly accurate BO, adiabatic, relativistic and QED effects but no provision for non-adiabatic effects, therefore represent an accurate characterisation of the value of non-adiabatic effects for each \\hp\\ level. Allowance for non-adiabatic effects using the simple model of PT \\cite{jt236} results in a further reduction of this deviation to 0.22~\\cm. Use of the non-adiabatic model of Diniz \\etal shows that in this model the use of QED corrections reduces the errors in the results by almost a factor of two from 0.33~\\cm\\ to 0.19~\\cm. This demonstrates the necessity of including QED corrections in  accurate {\\it ab initio} treatments of \\hp\\ rotation-vibration energy levels;  it opens the way for the development of an accurate non-adiabatic model which could potentially reach the $10^{-2}$ \\cm\\  accuracy necessary for the assignment of Carrington -- Kennedy \\cite{CB82} near-dissociation spectrum of \\hp\\ and its isotopologues."
    },
    "1402/1402.2350_arXiv.txt": {
        "abstract": "We used six simultaneous {\\em XMM-Newton} and {\\em Rossi X-ray Timing Explorer} plus five {\\em Suzaku} observations to study the continuum spectrum and the iron emission line in the neutron-star low-mass X-ray binary 4U 1636$-$53. We modelled the spectra with two thermal components (representing the accretion disc and boundary layer), a Comptonised component (representing a hot corona), and either a Gaussian or a relativistic line component to model an iron emission line at $\\sim 6.5$ keV. For the relativistic line component we used either the {\\sc diskline}, {\\sc laor} or {\\sc kyrline} model, the latter for three different values of the spin parameter. The fitting results for the continuum are consistent with the standard truncated disc scenario. We also find that the flux and equivalent width of the iron line first increase and then decrease as the flux of the Comptonised component increases. This could be explained either by changes in the ionisation state of the accretion disc where the line is produced by reflection, or by light bending of the emission from the Comptonised component if the height at which this component is produced changes with mass accretion rate. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) consist of a compact object (a neutron star or a black hole) and a late-type companion star with a mass of less than $\\sim $1 M$_{\\odot}$. Material from the outer layers of the companion is stripped off and accretes onto the compact object via the inner Lagrangian point and an accretion disc \\citep{accretion_book}. The inner parts of the accretion disc, and in the case of neutron stars the neutron-star surface and boundary layer, emit mostly in the X-ray band. These systems show also high-energy (up to a few 100 keV) emission, likely produced by Comptonisation of the soft X-ray photons in a hot electron corona \\citep{mcc2000,sanna13}. An extra hard tail has been observed in some LMXBs, e.g., GX 17$+$2, 4U 1636$-$53, GX 349$+$2, Sco X-1, 4U 1608$-$522, GX 13$+$1, 4U 1705$-$44 \\citep{salvo00,fiocchi06,paizis06,piraino07}.  Based on X-ray spectral and rapid variability properties, \\citet{hasinger89} classified the neutron-star low-mass X-ray binaries into two categories: the Z sources and the Atoll sources, owing their names to the shapes that the source traces in an X-ray colour-colour diagram. The Z sources have higher luminosities ($\\sim 0.5-1~L_{\\rm Edd}$) than the Atoll sources \\citep[$0.01-0.2 ~L_{\\rm Edd}$; e.g.,][]{done07,homan07,ford2000}. In Atoll sources the X-ray spectrum softens and the time-scale of the majority of the variability components decreases as the luminosity of the source generally increases \\citep{hasinger89,van03}. More specifically, at low luminosities these sources are in the so-called hard state, in which the Comptonised component dominates the energy spectrum; this component can be reasonably described by a power law with a photon index of $\\sim 1.6 - 2.5$ \\citep{yoshida93,mendez97}. The temperature at the inner edge of the accretion disc is relatively low, $0.3 - 0.5$ keV \\citep{sanna13}, and the disc usually contributes less than 20\\% of the emission in the $1-20$ keV energy range. In the truncated disc scenario \\citep[see, e.g.,][and references therein]{done07}, the accretion disc is truncated at a large radius in the hard state, that being the reason for the relatively low inner-disc temperature and thermal component flux \\citep{gierlinski03}. At high luminosities, in the soft state, the disc emission in the $1-20$ keV range becomes more significant. The standard accretion-disc model \\citep{ss73} predicts that in this case the disc extends down to the innermost stable circular orbit radius, leading to a high disc temperature and a strong thermal component. Compared to the hard state, in the soft state the inner-disc temperature increases to $0.7-1.0$ keV \\citep{sanna13}; since the number of soft photons increases as $T^4$, the electrons in the corona are efficiently cooled down via the inverse Compton scattering process, and the Comptonised spectrum steepens \\citep[photon index $\\sim 2 - 2.5$, e.g.,][]{miya93,mendez97}. In the soft state, the thermal components dominate the X-ray spectrum below $\\sim 20$ keV, and little hard emission is detected (Gierli\\'nski \\& Done. 2003). Both in the hard and the soft state the neutron star surface or boundary layer, usually fitted with a blackbody component \\citep{white88,sanna13}, contributes significantly to the emission at energies below $\\sim 10$ keV. Going from the hard to the soft state, or vice-versa, the spectra of these sources display some intermediate properties, and the source is said to be in the intermediate, or transitional, state. (These basic states receive several names depending on the class of the source; see, e.g., \\citealt{hasinger89}). The mechanism driving the transition between the hard to the soft state is still unclear. However, it is generally assumed that the change of mass accretion rate and the disc geometry are connected to the state evolution \\citep{hasinger89,mendez99,lin07}.  Besides the emission components described above, the accretion disc is likely illuminated by the Comptonised photons and the thermal spectrum from the neutron star and its boundary layer, and as a consequence it produces a reflection spectrum \\citep[e.g.][]{fabian10}. Due to the high abundance and fluorescence yield, an iron emission line at $\\sim$6$-$7 keV may appear in the spectrum, with an intrinsic line width of the order of 1eV \\citep{basko78}. The iron line profile is asymmetrically broadened by the fast disc rotation (Doppler-broadening) and special and general relativistic effects (e.g., Doppler boosting and gravitational redshift) near the central compact object \\citep{fabian89}. The final line profile is determined by parameters of the system, like the inclination angle of the disc with respect to the line of sight, the inner radius of the disc, and the spin parameter of the central object. Thus measurements of the iron line profile provide an excellent way to study the physics and geometry of the accretion process \\citep{bhatta07,cackett08,miller13}.  4U 1636$-$53 is an Atoll LMXBs, consisting of a neutron star and a 0.4 M$_{\\odot}$ companion in a 3.8 hr orbit \\citep{pedersen82}, at 6 kpc distance \\citep{galloway06}. The source shows the full range of spectral states \\citep{belloni07,alta08}. Highly coherent burst oscillations indicate that the system harbors a millisecond pulsar with a spin frequency of 581 Hz \\citep{zhang97,stro02}. A pair of quasi-periodic oscillations (QPOs) at kHz frequencies were discovered by \\citet{zhang96} and \\citet{wijnands97}. \\citet{kaaret99} and \\citet{marcio13} found that the soft X-ray emission in the lower kHz QPO (from the pair of kHz QPOs the one at lower frequency) lags the hard X-ray emission, suggesting that the emission is due to reprocessing of hard X-rays in a cooler Comptonising corona or the accretion disc, in a region with a size of at most a few kilometres. In the last ten years 4U 1636--53 showed a regular state transition cycle of $\\sim$40 days \\citep{shih05,belloni07}, making it an excellent source to study the variations of the broadband spectrum and iron line as a function of spectral state.  In this work, we study the spectrum of 4U 1636$-$53 in different states. We use eleven observations from Suzaku, XMM-Newton, and the Rossi X-Ray Timing Explorer (RXTE) satellites covering a wide range in luminosity, allowing us to investigate the evolution of the different spectral components as a function of spectral state, and study possible correlations between the continuum spectrum and the iron line. Since the disc is illuminated and photon-ionised by the continuum emission to produce the iron line, correlations between the iron line flux and the flux of different continuum components may provide an important clue to understand the origin of the iron emission line and the evolution of the accretion flow geometry. ",
        "conclusions": "We used all available Suzaku and XMM-Newton observations, the latter complemented with simultaneous RXTE observations, to study the evolution of the continuum spectrum and the iron line emission in the neutron star low-mass X-ray binary 4U 1636$-$53 across different spectral states. We found that the temperature of the neutron star and that at the inner edge of the accretion disc increase, whereas the electron temperature of the corona decreases, as the source moves across the colour-colour diagram from the hard to the soft state. Simultaneously, the inner radius of the accretion disc deduced from the {\\sc diskbb} component decreases rapidly by a factor of $\\sim 10$, and then remains more or less constant as the source reaches the soft state, suggesting that the inner edge of the disc may have reached the ISCO. The power-law photon index of the component used to fit the corona emission increases from $\\sim 1.7$ in the hard state up to $\\sim 2.8$ (depending on the model used to fit the line) in the soft state. Simultaneously, the total unabsorbed flux and the flux of the Comptonised component remain more or less constant, whereas the flux of the disc and that of the neutron star plus boundary layer increase, as the source moves from across the colour-colour diagram from the hard to the soft state (all fluxes calculated in the $0.5 -130$ keV range). Interestingly, the flux and equivalent width of the iron line first increase and then decrease as the flux of the Comptonised component increases. The maximum of this relation takes place when the source is close to the vertex in the colour-colour diagram.  The relation between the flux or the equivalent width of the line on one side and the flux of the hard spectral component on the other appears to contradict the expected behaviour if the line is due to reflection off the accretion disc. In the reflection scenario, the corona (and in neutron-star systems possibly also the neutron-star or boundary-layer) illuminates and photon-ionises the disc, where the reflection spectrum (continuum + emission line) is then produced. In this scenario the flux or equivalent width of the iron line should be positively correlated with that of the illuminating component. We find that, while this is the case when at low {\\sc nthcomp} flux values, the opposite is true when the {\\sc nthcomp} flux is high (see Figure \\ref{corrflx}). However, the disc becomes more ionised when the illuminating flux increases, provided that there are enough photons with energies above the ionisation potential of H- and He-like iron. As the flux of the illuminating component increases further, the material in the disc would eventually become fully ionised, and the line will disappear. This scenario agrees with recent calculations of the reflection spectrum of an ionised slab at different ionisation levels  \\citep{garcia13}. \\cite{garcia13} found (see their Figure 5) that the flux of the iron emission line in the reflected spectrum decreases gradually as the material in the slab is ionised further. Similarly, \\cite{matt93} showed that the equivalent width of the iron line in the reflected spectrum of an ionised slab initially increases and then decreases as the flux of the ionising source increases. Our results are generally consistent with this idea \\citep[compare our Figure \\ref{corrflx} with  Figures 1 and 2 in][]{matt93}. However, in this scenario, as the ionisation state of the disc increases, the fitted energy of the emission line should gradually increase, as the fraction of H-like iron ions relative to the He-like iron ions in the disc increases. Our results do not show any correlation between the fitted energy of the line and the flux of the ionising source (see Tables \\ref{sline} and \\ref{xline}), regardless of the model we used to fit the line. While this lack of correlation could be due to limitations of the models that we used to fit the iron line, or to the fact that we only fitted the iron line instead of the full reflection spectrum, using the same XMM-Newton data presented here \\citet{sanna13} found no correlation between the ionisation parameter of the full reflection model they fitted to the spectra and the position of the source in the colour-colour diagram, reinforcing this conclusion.  Alternatively, the relation between iron line flux or equivalent width and the flux of the hard spectral component in 4U 1636--53 shown in Figure \\ref{corrflx} could be interpreted in terms of light bending \\citep{miniutti04}. In this model, the observed flux of the hard illuminating source (in our fits the {\\sc nthcomp}) depends strongly upon the height of the source of photons in the corona above the disc; the observed flux of the direct emission can change by up to a factor $\\sim20$ as the height of the illuminating source changes, even if the luminosity of the illuminating source remains constant. On the other hand, the reflected spectrum of the disc, and hence the flux or equivalent width of the iron line, is much less sensitive to the height of the illuminating source. \\cite{miniutti04} find that there are three regimes in which, as the flux of the direct emission increases, the reflection component (and hence the iron line flux or equivalent width) is first correlated, then insensitive, and finally anti-correlated with the flux of the direct emission. \\cite{rossi05} found that, in the black-hole candidate XTE J1650--500, as the total flux of the source above 7 keV increases, the flux of the iron line initially remains constant and eventually decreases. They found that the source flux level at which the line flux starts to decrease coincides with the transition from the hard to the soft state in this source. \\cite{park04} also found a complex relation between the iron-line flux and the flux of the hard and the soft components in the black-hole candidate 4U 1543--475. In this case, however, the relation does not follow any clear trend, and an interpretation in terms of the light bending model is not apparent. Rossi et al. (2005) suggested the direct and reflected components in XTE J1650--500 could be related to the existence of a radio jet in this source \\citep{corbel04}, with possible changes of the height of the base of the jet as the mechanism that drives the changes of the direct emission in XTE J1650--500. It is unclear whether this could also be the case in 4U 1636--53, since this source has not been detected in radio, and hence the existence of a strong jet in this system is at best doubtful \\citep{thomas79,russell12}. The lack of a strong radio jet in 4U 1636--53 could also be the reason that the flux of the hard (direct) component in this source changes only by a factor $\\sim 5$, whereas in XTE J1650--500 the flux of the direct ionising source changes by $\\simmore 10$. We finally note that if the relation between the flux or equivalent width of the iron line and the {\\sc nthcomp} flux in 4U 1636--53 is due to light bending, this would be the first case of a neutron-star system in which this effect is observed. The light bending model was developed for the case of a rapidly spinning black hole, and therefore it is unclear whether it would also apply for a moderately spinning neutron star.  In the case of 4U 1636--53 the point at which the flux of the iron line switches from correlated to anti-correlated with the flux of the {\\sc nthcomp} component coincides with the vertex in the colour-colour-diagram, at $S_{a} \\sim 2.1$. The existence of the vertex in this diagram directly indicates a sudden change of the spectral properties of the source which, going from the hard to the soft state, quickly softens there. We note also that the vertex of the colour-colour diagram is the place where the quality factor of the kHz QPOs \\citep{barret06,mendez06} and the intrinsic coherence between the variability in the hard and soft energy bands \\citep{marcio13} are the highest. Furthermore, \\citet{zhang11} studied 298 type-I X-ray bursts in the 4U 1636$-$53 using RXTE observations, and they found that in this area of the colour-colour diagram most photospheric radius expansion bursts and a super-burst are observed. All these properties suggest that, in this source several properties of the accretion flow change significantly in this area of the colour-colour diagram.  The changes of the properties of the continuum spectrum of 4U 1636--53 as the source moves in the colour-colour diagram are generally consistent with those of other neutron-star LMXBs. For instance, \\citet{farinelli11} measured the power-law index and the electron temperature $kT_{\\rm e}$ of the Comptonising component in the neutron star LMXBs Sco X-1, GX 349$+$2, X 1658$-$298, 1E 1724$-$3045, GX 17$+$2, Cyg X$-$2, GX 340$+$0, GX 3$+$1, and GS 1826$-$238, and found that the power-law energy index ($\\Gamma -1$ in the {\\sc nthcomp} component) remains more or less constant at around 1$\\pm$0.2. Using RXTE and BeppoSAX observations, \\citet{seifina11,seifina12} showed that in 4U 1728$-$34 and GX 3$+$1, the power-law index also remains almost constant as the temperature of the corona changes dramatically. Recently, \\citet{titar13} found that in another neutron star binary, 4U 1820$-$30, the power-law index remains almost constant for different source states. The power-law index in 4U 1636$-$53 changes from $\\sim 1.7$ in the hard state to $\\sim 2.8$ in the soft state (depending on the model used to fit the iron line), which is a somewhat larger than the variations in the sources studied by \\citet[][note, however, that the model that they used to fit the Comptonised component is not the same as the one we used here]{farinelli11}. \\citet{titar13} concluded that the power-law index quasi-stability is an intrinsic property of neutron star binaries, which is fundamentally different from that of black-hole binary systems. \\citet{seifina11} suggested that this stability of the power-law index happens when the energy released in the corona itself is much higher than the one from the disc intercepted by the corona.  The changes of the continuum spectrum are also broadly consistent with the truncated disc scenario in LMXBs \\citep[e.g.][and references therein]{done07}: The temperature of the disc increases and the inner radius of the disc decreases as the inferred mass accretion rate increases. The electron temperature of the {\\sc nthcomp} component and the temperature of the neutron-star surface are also consistent with this scenario. It is intriguing, however, that the inner radius of the accretion disc deduced from the relativistically broaden iron line does not follow the same trend. The model of the disc that we used to fit the data does not include spectral hardening, and the iron line profile may be further affected by mechanisms other than relativistic broadening. It remains to be seen whether these effect could explain this discrepancy."
    },
    "1402/1402.4483_arXiv.txt": {
        "abstract": "We report on a comprehensive X-ray spectral analysis of the nearby radio-quiet quasar MR~2251$-$178, based on the long-look ($\\sim 400$ ks) \\textit{XMM--Newton} observation carried out in November 2011. As the properties of the multiphase warm absorber (thoroughly discussed in a recent, complementary work) hint at a steep photoionizing continuum, here we investigate into the nature of the intrinsic X-ray emission of MR~2251$-$178 by testing several physical models. The apparent 2--10 keV flatness as well as the subtle broadband curvature can be ascribed to partial covering of the X-ray source by a cold, clumpy absorption system with column densities ranging from a fraction to several $\\times 10^{23}$ cm$^{-2}$. As opposed to more complex configurations, only one cloud is required along the line of sight in the presence of a soft X-ray excess, possibly arising as Comptonized disc emission in the accretion disc atmosphere. On statistical grounds, even reflection with standard efficiency off the surface of the inner disc cannot be ruled out, although this tentatively overpredicts the observed $\\sim 14$--150 keV emission. It is thus possible that each of the examined physical processes is relevant to a certain degree, and hence only a combination of high-quality, simultaneous broadband spectral coverage and multi-epoch monitoring of X-ray spectral variability could help disentangling the different contributions. Yet, regardless of the model adopted, we infer for MR~2251$-$178 a bolometric luminosity of $\\sim 5$--$7 \\times 10^{45}$ erg s$^{-1}$, implying that the central black hole is accreting at $\\sim15$--25 per cent of the Eddington limit. ",
        "introduction": "Among the X-ray brightest Active Galactic Nuclei (AGN) in the local Universe due to its 2--10 keV luminosity largely exceeding 10$^{44}$ erg s$^{-1}$, MR~2251$-$178 ($z \\simeq 0.064$; Canizares, McClintock \\& Ricker 1978) is a spectacular object in every respect. It was the first quasar to be detected and identified through X-ray observations (Cooke et al. 1978; Ricker et al. 1978), as well as the first one where the presence was established of \\textit{warm} absorption by photoionized gas, variable in both ionization state and possibly column density over timescales of less than one year (Halpern 1984). The source was later found to experience appreciable changes in the X-ray flux over periods of $\\sim 10$ days (Pan, Stewart \\& Pounds 1990), with a tight correlation between the ionization parameter of the absorbing gas\\footnote{The ionization parameter is defined as $\\xi = L_\\rmn{ion}/nr^2$, where $n$ is the electron density of the gas and $r$ is its distance from an ionizing source with 1--1000 Ry luminosity $L_\\rmn{ion}$ (Tarter, Tucker \\& Salpeter 1969).} and the continuum luminosity (Mineo \\& Stewart 1993). Narrow absorption lines with a systematic blueshift of $\\sim 300$ km s$^{-1}$ have been detected in the ultraviolet (UV) due to Ly$\\alpha$, N~\\textsc{v} and C~\\textsc{iv} (Monier et al. 2001). In particular, the C~\\textsc{iv} doublet apparently vanished in less than four years (Ganguly, Charlton \\& Eracleous 2001), implying that the UV absorber is truly local to the AGN, and that it is possibly one and the same with the soft X-ray warm absorber. Much deeper insights came with the advent of high-resolution X-ray spectroscopy. Early \\textit{XMM--Newton} Reflection Grating Spectrometer (RGS) observations hinted at a multiphase configuration for the warm absorber, likely consisting of two or three distinct components (Kaspi et al. 2004), while the \\textit{Chandra} High Energy Transmission Grating (HETG) spectrum revealed a highly ionized absorption feature in the iron K band, interpreted as the Fe~\\textsc{xxvi} Ly$\\alpha$ line at the sizable outflow velocity of $v_\\rmn{out} \\sim 0.04c$ (Gibson et al. 2005). The corresponding mass-loss rate was calculated to be at least an order of magnitude larger than the accretion rate, unless the covering fraction of the outflow is very small.  The environment of MR~2251$-$178 is likewise exceptional. The quasar lies in the outskirts of a loose, irregular cluster with several tens of galaxies (Phillips 1980), and is surrounded by a huge emission-line nebula detected in H$\\alpha$ and [O~\\textsc{iii}] at optical wavelengths out to a distance of $\\sim 100$ kpc from the central source (Bergeron et al. 1983; Shopbell, Veilleux \\& Bland-Hawthorn 1999). The knotty and filamentary appearance of this gaseous envelope is more typical of powerful radio galaxies, yet MR~2251$-$178 only shows weak radio emission, whose elongated morphology resembles a double-lobed jet-like structure (Macchetto et al. 1990). A similar spatial extent and orientation marks the ionization cones recently found through deep [O~\\textsc{iii}]$\\lambda$5007/H$\\beta$ and [N~\\textsc{ii}]$\\lambda$6583/H$\\alpha$ flux ratio maps (Kreimeyer \\& Veilleux 2013). Overall, the quasar radiation field can easily account for the ionization state of the nebula, which is arguably the most extended around a radio-quiet source. Its origin, however, is still unclear. Diffuse X-ray emission along some directions was preliminary reported by Gibson et al. (2005) based on the smoothed HETG zeroth-order image taken in 2002. Unfortunately, we cannot safely corroborate these findings through the much deeper 2011 HETG data set, since the artificial broadening of the instrumental Point Spread Function due to photon pile-up overrides any faint contribution from the halo.\\footnote{The AGN flux in the 2011 \\textit{Chandra} observation was larger by $\\sim 50$ per cent with respect to 2002. While the effects of pile-up on the Point Spread Function can be qualitatively simulated, they prevent any accurate image reconstruction.}  Due to its key role as the underlying source of ionizing photons for the ambient gas, from nuclear to intergalactic scales, a proper knowledge of the shape and behaviour of the intrinsic X-ray continuum is highly desirable, yet most of the effort in the wealth of X-ray analyses performed so far has concentrated on improving the characterization of the multi-component warm absorber. Incidentally, all the X-ray spectral studies actually agree with describing the absorbed continuum by means of a power law of photon index $\\Gamma \\sim 1.6$, with a high-energy exponential cutoff at $\\sim 100$ keV (Orr et al. 2001) and a soft excess below $\\sim 1$ keV (Kaspi et al. 2004). In a recent paper, we have taken advantage of coordinated $\\sim 400+400$ ks long \\textit{Chandra} HETG and \\textit{XMM--Newton} RGS observations to bring our grasp on the properties of the intervening ionized gas to unprecedented detail and energy resolution (Reeves et al. 2013; hereafter R13). Both campaigns were conducted as large programs in late 2011, with just a few weeks of separation from one another, and yielded intriguing information on the illuminating continuum itself, which below 2 keV is required to be much steeper ($\\Gamma > 2$) than typically estimated at hard X-rays. Indeed, the analysis of the \\textit{Suzaku} spectrum carried out by Gofford et al. (2011) had already suggested that the broadband X-ray emission of MR~2251$-$178 can be reproduced equally well through a softer power law with $\\Gamma \\simeq 2$, provided that an additional cold-gas column of $N_\\rmn {H} \\sim 10^{23}$ cm$^{-2}$ covering a moderate fraction of the source is introduced.  In the wake of these indications, here we present a thorough investigation of the 0.3--10 keV EPIC/pn spectrum obtained in the 2011 \\textit{XMM--Newton} observation, with the aim of understanding the nature of the primary photoionizing continuum and unveiling the high-energy physical processes at work in the very central regions of this powerful quasar. This work is organized as follows. In Section~2 we provide the basic details about the \\textit{XMM--Newton} large observing program on MR~2251$-$178, and describe the principal steps of our data reduction. Section~3 is dedicated to the spectral analysis, whose results and main implications are discussed in Section~4. Our conclusions are drawn in Section~5. Throughout this paper we have assumed $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$ and $\\Omega_\\Lambda=0.73$, in agreement with the latest values of the concordance cosmological parameters (Hinshaw et al. 2013). ",
        "conclusions": "Based on a recent \\textit{XMM--Newton} long-look observation with $\\sim 270$ ks of net exposure, we have reported on the 0.3--10 keV EPIC/pn spectral analysis of MR~2251$-$178, one of the brightest radio-quiet quasars in the local Universe. Following a companion paper dedicated to the study of the high-resolution grating spectra and of the physical properties of the complex warm absorber, here we have focused on the nature of the intrinsic X-ray emission. The broadband X-ray continuum of MR~2251$-$178 is known to exhibit substantial curvature up to $\\sim 100$ keV, where a possible cutoff is also present. Together with the apparent soft excess below 0.7 keV and the warm absorption trough at $\\sim 1$--2 keV, this gives rise to a peculiar steep--flat--steep spectral shape, which has been inspected within the frameworks of partial covering absorption, ionized reflection, and Comptonized disc emission.  As their application turns out to be nearly statistically equivalent over the 0.3--10 keV band, all these models have also been compared with the coeval OM photometric data and the $\\sim 14$--150 keV \\textit{Swift}/BAT spectrum, averaged over 70 months. Assuming no cross normalization, the entire SED shows that the optical and hard X-ray fluxes lie at almost the same level in $\\nu f_\\nu$ units. Indeed, the optical to X-ray spectral index is extremely small ($\\alpha_\\rmn{ox} < 1.2$), irrespective of the exact absorption correction. The corresponding bolometric output of MR~2251$-$178 is estimated to be $L_\\rmn{bol} \\sim 5$--$7 \\times 10^{45}$ erg s$^{-1}$, i.e. $\\sim 15$--25 per cent of the source Eddington luminosity.  None of the alternative interpretations taken into account in this work are conclusive by themselves, displaying both strengths and limitations at the same time, which can be summarized as follows:  1) The pure partial covering scenario reveals two distinct low-ionization components with columns of 0.6 and $6.6 \\times 10^{23}$ cm$^{-2}$, eclipsing a fraction of $\\sim 0.23$ and 0.49 of the X-ray source, respectively. We argue that these absorbers are located at BLR scales, but are unlikely to belong to different layers. Density gradients might be entailed, pointing to a single clump of gas with a dense core and a lighter halo. The intrinsic photon index $\\Gamma \\sim 2.2$ is close to what is required by the photoionization models of the warm absorber, but still somewhat lower; even so, excess curvature is left at hard X-rays around 50 keV.  2) At EPIC/pn resolution, X-ray reflection provides a reasonable fit without implying any extreme gravity regime. The blurred disc component, however, is unusually cold, and $\\sim 50$ times less ionized than its distant counterpart. This is possibly an artefact to compensate for the inadequate results with the soft X-ray emission lines (resolved by the gratings), and also leads to a misrepresentation of the warm absorption properties (for which the illumination is anyway flat). The degree of smoothing appears too large for the iron K band, and a clear pattern is found in the residuals above 9 keV. Moreover, the BAT spectrum is largely overpredicted.  3) The optically-thick/thin Comptonization of thermal photons from the inner disc into both the soft excess and the hard power law, with the outer regions emitting as a colour temperature corrected blackbody, is established as the most obvious connection between the optical and X-ray bands. Nevertheless, in spite of the remarkably good extrapolation to the OM data and of the ample leverage on the physical parameters, a basic self-consistent model either fails to reproduce the 0.3--10 keV features or invariably plummets too quickly at higher energies. Any agreement is lost on the three warm absorption components.  A composite model allows us to successfully overcome most of the above shortcomings. While reflection from the accretion flow and/or the ambient material might still be involved to a certain extent, in the preferred picture the hard power-law continuum ($\\Gamma \\sim 1.75$) steepens below $\\sim 1$ keV into a soft excess with $kT \\sim 0.14$ keV, whose origin remains unclear although it can be tentatively associated with Comptonization in the disc atmosphere. A single cloud with $N_\\rmn{H} \\sim 3 \\times 10^{23}$ cm$^{-2}$ is present along the line of sight, covering about one sixth of the X-ray source.  In conclusion, the X-ray observations of MR~2251$-$178 brought to light a complex and stratified environment close to the central source, which deeply transforms the shape of the intrinsic X-ray spectrum. Our study suggests that only a combination of time variability analysis, simultaneous broadband coverage and high spectral resolution can help disentangling the effects of the different physical processes responsible for the observed X-ray emission of AGN."
    },
    "1402/1402.3677.txt": {
        "abstract": "The emission of radio waves from air showers has been attributed to the so-called geomagnetic emission process. At frequencies around 50 MHz this process leads to coherent radiation which can be observed with rather simple setups. The direction of the electric field induced by this emission process depends only on the local magnetic field vector and on the incoming direction of the air shower. We report on measurements of the electric field vector where, in addition to this geomagnetic component, another component has been observed which cannot be described by the geomagnetic emission process. The data provide strong evidence that the other electric field component is polarized radially with respect to the shower axis, in agreement with predictions made by Askaryan who described radio emission from particle showers due to a negative charge-excess in the front of the shower. Our results are compared to calculations which include the radiation mechanism induced by this charge-excess process. ",
        "introduction": "\\label{sec:introduction}  When high-energy cosmic rays penetrate the atmosphere of the Earth they induce an air shower. The detailed registration of this avalanche of secondary particles is an essential tool to infer properties of the primary cosmic ray, such as its energy, its incoming direction, and its composition. Radio detection of air showers started in the 1960's and the achievements in these days have been presented in reviews by Allan \\cite{Refworks:188} and Fegan \\cite{Refworks:189}. In the last decade, there has been renewed interest through the publications of the LOPES \\cite{Refworks:101} and CODALEMA \\cite{Refworks:91} collaborations. We have deployed and are still extending the Auger Engineering Radio Array (AERA) \\cite{Refworks:178,Refworks:180,Refworks:183,vandenBergECR2012,SchroederICRC2013} as an additional tool at the Pierre Auger Observatory to study air showers with an energy larger than $10^{17}$ eV. In combination with the data retrieved from the surface-based particle detectors \\cite{Refworks:176} and the fluorescence detectors \\cite{Refworks:177} of this observatory, the data from radio detectors can provide additional information on the development of air showers.  %An important step in the interpretation of the data obtained with radio-detection methods is the understanding of the emission mechanisms. In the early studies of radio emission from air showers it was conjectured that two emission mechanisms play an important role: the geomagnetic emission mechanism as proposed, amongst others, by Kahn and Lerche \\cite{Refworks:190} and the charge-excess mechanism as proposed by Askaryan \\cite{Refworks:191}.  %BEGIN NEW TEXT  An important step in the interpretation of the data obtained with radio-detection methods is the understanding of the emission mechanisms. In the early studies of radio emission from air showers it was conjectured that two emission mechanisms play an important role: the geomagnetic emission mechanism as proposed, amongst others, by Kahn and Lerche \\cite{Refworks:190} and the charge-excess mechanism as proposed by Askaryan \\cite{Refworks:191}. Essential for the geomagnetic effect is the induction of a transverse electric current in the shower front by the geomagnetic field of the Earth while the charge excess in the shower front is to a large extent due to the knock-out of fast electrons from the ambient air molecules by high-energy photons in the shower. The magnitude of the induced electric current as well as the induced charge excess is roughly proportional to the number of particles in the shower and thus changing in time. The latter results in the emission of coherent radio waves at sufficiently large wavelengths \\cite{Refworks:190, REF12a}. The shower front, where both the induced transverse current and the charge excess reside, moves through the air with nearly the velocity of light. Because air has a refractive index which differs from unity, Cherenkov-like time compression occurs \\cite{REF12b, REF12c}, which affects both the radiation induced by the transverse current as well as by the charge excess. The polarization of the emitted radiation differs for current-induced and charge-induced radiation, but its direction for each of these individual components does not depend on the Cherenkov-like time compression caused by the refractive index of air. For this reason we will distinguish in this paper only geomagnetic (current induced) and charge-excess (charge induced) radiation. The contribution of the geomagnetic emission mechanism, described as a time-changing transverse current by Kahn and Lerche \\cite{Refworks:190}, has been observed and described in several papers; see e.g. Refs. \\cite{Refworks:101,Refworks:74,Refworks:35,Refworks:186,Refworks:221}. Studies on possible contributions of other emission mechanisms from air showers have also been reported \\cite{Refworks:208,Refworks:209,Refworks:210,Refworks:211,Prescott1971}. An observation of the charge-excess effect in air showers has been reported by the CODALEMA collaboration \\cite{Refworks:198}.   %END NEW TEXT  We present the analysis of two data sets obtained with two different setups consisting of radio-detection stations (RDSs) deployed at the Pierre Auger Observatory. The first data set was obtained with a prototype setup \\cite{Refworks:67} for AERA; the other one with AERA itself \\cite{Refworks:178,Refworks:180,Refworks:183} during its commissioning phase while it consisted of only 24 stations. In addition, we will compare these data sets with results from different types of calculations outlined in Refs. \\cite{Refworks:3, Refworks:11, Refworks:196, Refworks:220, Refworks:212, Refworks:213, Refworks:223}.  This paper is organized as follows. We discuss in Section \\ref{sec:detectionsystem} the experimental equipment used to collect our data. In Section \\ref{sec:dataanalysis} we present the data analysis techniques and the cuts that we applied on the data, whilst in Section \\ref{sec:comparison} we compare our data with calculations. In Section \\ref{sec:conclusion} we discuss the results and we present our conclusions. For clarity, Sections \\ref{sec:dataanalysis} and  \\ref{sec:comparison} contain only those figures which are based on the analysis for the data obtained with AERA during the commissioning of its first 24 stations; the results of the prototype are shown in Appendix C. We mention that analyses of parts of the data have been presented elsewhere \\cite{Refworks:186,Refworks:199,Refworks:214,TimHuegeICRC2013}.  %\\clearpage ",
        "conclusions": "\\label{sec:conclusion}  We have studied with two different radio-detection setups deployed at the Pierre Auger Observatory the emission around 50 MHz of radio waves from air showers. For a sample of 37 air showers the electric field strength has been analyzed as a tool to disentangle the emission mechanism caused by the geomagnetic and the charge-excess processes. For the present data sets, the emission is dominated by the geomagnetic emission process while, in addition, a significant fraction of on average $(14 \\pm 2)\\%$ is attributed to a radial component which is consistent with the charge-excess emission mechanism. Detailed simulations have been performed where both emission processes were included. The comparison of these simulations with the data underlines the importance of including the charge-excess mechanism in the description of the measured data. However, further refinements of the models might be required to fully describe the present data set. A possible reason for the incomplete description of the data by the models might be an underestimate of (systematic) errors in the data sets or the effect of strong electric fields in the atmosphere.   The successful installation, commissioning, and operation of the Pierre Auger Observatory would not have been possible without the strong commitment and effort from the technical and administrative staff in Malarg\\\"ue.  We are very grateful to the following agencies and organizations for financial support: Comisi\\'on Nacional de Energ\\'ia At\\'omica, Fundaci\\'on Antorchas, Gobierno De La Provincia de Mendoza, Municipalidad de Malarg\\\"ue, NDM Holdings and Valle Las Le\\~nas, in gratitude for their continuing cooperation over land access, Argentina; the Australian Research Council; Conselho Nacional de Desenvolvimento Cient\\'ifico e Tecnol\\'ogico (CNPq), Financiadora de Estudos e Projetos (FINEP), Funda\\c{c}\\~ao de Amparo \\`a Pesquisa do Estado de Rio de Janeiro (FAPERJ), S\\~ao Paulo Research Foundation (FAPESP) Grants \\#2010/07359-6, \\#1999/05404-3, Minist\\'erio de Ci\\^{e}ncia e Tecnologia (MCT), Brazil; AVCR, MSMT-CR LG13007, 7AMB12AR013, MSM0021620859, and TACR TA01010517 , Czech Republic; Centre de Calcul IN2P3/CNRS, Centre National de la Recherche Scientifique (CNRS), Conseil R\\'egional Ile-de-France, D\\'epartement  Physique Nucl\\'eaire et Corpusculaire (PNC-IN2P3/CNRS), D\\'epartement Sciences de l'Univers (SDU-INSU/CNRS), France; Bundesministerium f\\\"ur Bildung und Forschung (BMBF), Deutsche Forschungsgemeinschaft (DFG), Finanzministerium Baden-W\\\"urttemberg, Helmholtz-Gemeinschaft Deutscher Forschungszentren (HGF), Ministerium f\\\"ur Wissenschaft und Forschung, Nordrhein-Westfalen, Ministerium f\\\"ur Wissenschaft, Forschung und Kunst, Baden-W\\\"urttemberg, Germany; Istituto Nazionale di Fisica Nucleare (INFN), Ministero dell'Istruzione, dell'Universit\\`a e della Ricerca (MIUR), Gran Sasso Center for Astroparticle Physics (CFA), CETEMPS Center of Excellence, Italy; Consejo Nacional de Ciencia y Tecnolog\\'ia (CONACYT), Mexico; Ministerie van Onderwijs, Cultuur en Wetenschap, Nederlandse Organisatie voor Wetenschappelijk Onderzoek (NWO), Stichting voor Fundamenteel Onderzoek der Materie (FOM), Netherlands; Ministry of Science and Higher Education, Grant Nos. N N202 200239 and N N202 207238, The National Centre for Research and Development Grant No ERA-NET-ASPERA/02/11, Poland; Portuguese national funds and FEDER funds within COMPETE - Programa Operacional Factores de Competitividade through Funda\\c{c}\\~ao para a Ci\\^{e}ncia e a Tecnologia, Portugal; Romanian Authority for Scientific Research ANCS, CNDI-UEFISCDI partnership projects nr.20/2012 and nr.194/2012, project nr.1/ASPERA2/2012 ERA-NET, PN-II-RU-PD-2011-3-0145-17, and PN-II-RU-PD-2011-3-0062, Romania; Ministry for Higher Education, Science, and Technology, Slovenian Research Agency, Slovenia; Comunidad de Madrid, FEDER funds, Ministerio de Ciencia e Innovaci\\'on and Consolider-Ingenio 2010 (CPAN), Xunta de Galicia, Spain; The Leverhulme Foundation, Science and Technology Facilities Council, United Kingdom; Department of Energy, Contract Nos. DE-AC02-07CH11359, DE-FR02-04ER41300, DE-FG02-99ER41107, National Science Foundation, Grant No. 0450696, The Grainger Foundation USA; NAFOSTED, Vietnam; Marie Curie-IRSES/EPLANET, European Particle Physics Latin American Network, European Union 7th Framework Program, Grant No. PIRSES-2009-GA-246806; and UNESCO.  \\clearpage \\appendix"
    },
    "1402/1402.3111.txt": {
        "abstract": "{} %aims {Period study of 321 fundamental mode RR Lyrae type stars (RRab), which had appropriate data in ASAS and SuperWASP surveys, was performed to complement and extend the list of known Blazhko stars in galactic field with bright stars up to 12.5 mag in maximum light.} %methods {An individual approach was applied to each studied star. Permanent visual supervision was maintained to each procedure in data analysis (data cleaning, frequency spectra examination) to avoid missing any possible sign of the Blazhko effect. Period analysis was performed using \\textsc{Period04} software.} %results {We found 100 stars to be definitely modulated. In 25 cases, previously unknown modulation was revealed and 8 new candidates for Blazhko stars were identified. Their modulation needs to be confirmed. In 18 previously found Blazhko stars, no modulation was detectable. Multiple modulation was revealed for eight stars that were previously proposed to show simple modulation. In total, there were twelve stars with some peculiarity in their modulation in the sample. This brings the incidence rate of multiple/irregularly-modulated stars to 12~\\%. The ratio of the modulation periods of five of the double-modulated stars was within the ratios of small integers. One of stars studied, IK Hya, showed a very interesting frequency spectrum, which we interpret as changing Blazhko period between 71.81 and 75.57~days and an additional 1403-day-long cycle analogous to a four-year cycle of the prototype RR Lyr. The limits of the shorter period produce a beating period that is approximately twice as big as a 1403-day period. The newly revealed Blazhko star RZ~CVn seems to undergo changes in the amplitude of the modulation, as well as in the basic pulsation and Blazhko periods. We found that the incidence rate of the Blazhko RR Lyraes is at least 31~\\%, more likely even higher. It was also found that the majority of the Blazhko variables show triplet structures in their frequency spectra and that in 89~\\% of these cases, the peak with larger amplitude is on the right-hand side of the main pulsation component.} {} ",
        "introduction": "Data from automatic sky surveys provide an invaluable opportunity for finding new variables and for long-term monitoring of the sky. The most extensive survey of Galactic field variable star research are without doubt Polish ASAS \\citep[\\textit{All Sky Automated Survey}; see e.g.][]{pojmanski1997,pojmanski2002} and SuperWASP \\citep[\\textit{Super Wide Angle Search for Planets}, hereafter SW; ][]{pollaco2006,butters2010} managed by the University of Leicester. Both projects have been run for several years.  Field RR Lyraes have been studied many times on the basis of various sky surveys. Among others, \\citet{kinemuchi2006} and \\citet{wils2006} utilized data from \\textit{Northern sky variability survey} \\citep[NSVS, ][]{wozniak2004}, e.g. \\citet{kovacs2005}, \\citet{szczygiel2009} studied RR~Lyraes based on ASAS data, and \\citet{drake2013} dealt with RR Lyraes measured in \\textit{Catalina sky survey}\\footnote {http://www.lpl.arizona.edu/css}. Probably the most precise ground-based survey data were provided by \\textit {the Konkoly Blazhko survey}, which have been operated since 2004 \\citep[e.g.][]{jurcsik2009a,sodor2007,sodor2012}.  Crucial advances in Blazhko effect research have been made recently by space missions \\textit{CoRoT} and \\textit{Kepler}. Ultra-precise data provided by these space telescopes allowed findings of, say, cycle-to-cycle variations now known as period doubling \\citep{szabo2010,kolenberg2010,benko2010} or long-term changes in the Blazhko effect and excitation of additional modes \\citep{poretti2010,chadid2010,guggenberger2012}.  Large samples of Blazhko RR Lyrae stars in galactic bulge were also studied in the framework of MACHO and OGLE surveys \\citep{soszynski2011, moskalik2003}. Since the stars in the bulge are faint and are not included in ASAS or SW surveys, we did not deal with them.  Special papers based on ASAS data, which are at least partially devoted to Blazhko stars in our Galaxy, were published by \\citet{kovacs2005}, \\citet{wils2005} (hereafter WS05), and \\citet{szczygiel2007} (hereafter SF07), who all revealed many new Blazhko variables and gave a list of them. An extensive study of the Blazhko effect is presented in the paper of \\citet{leborgne2012}, which is based on observations made with \\textit{TAROT} telescopes \\citep{klotz2008,klotz2009}. \\citet{leborgne2012} also made use of $O-C$ information stored in the \\textit{GEOS RR Lyr database} \\citep{leborgne2007} and gave a list of galactic field Blazhko stars. A regularly updated database of known Blazhko RR Lyraes situated in the galactic field \\citep[\\textit{BlaSGalF}, ][]{skarka2013b}\\footnote{http://physics.muni.cz/$\\sim$blasgalf/} recently contained about 270 objects.  There are a few studies of particular RR Lyraes proceeded from SW mesurements, e.g. the TV Boo study made by \\citet{skarka2013a}, or \\citet{srdoc2012}, who revealed the Blazhko effect in GSC02626-00896, but a comprehensive study of RR Lyraes based on SW data has not been performed yet.  RR Lyraes with appropriate ASAS and SW data were in the spotlight. The goal was to search for the presence of the Blazhko modulation and to describe its characteristics (pulsation and modulation periods and multiplicity/irregularity of the modulation). We applied commonly used techniques of light curve and frequency spectra examination to complement and extend the list of known Blazhko stars with bright RR Lyraes (up to 12.5 mag in maximum light). To be as precise as possible, we carefully cleaned and analysed each light curve and frequency spectra individually. ",
        "conclusions": "A period study of 321 fundamental mode galactic field RR Lyraes with brightness at maximum light higher than 12.5 mag was performed based on available data from the ASAS and/or SW database. These stars were chosen from a more extended sample consisting of 557 stars after careful pre-selection based on conditions described in sec. 2. Each light curve was cleaned and analysed individually. In addition, each frequency spectrum was scanned through careful visual inspection to avoid omitting any sign of the Blahzko effect.  One hundred stars were identified and confirmed to be Blazhko variables (25 previously unknown). Eight stars, marked as Blazhko candidates, showed some ambiguous indications of modulation like characteristic scattering around maximum of phased light curve. In contrast no indication of the Blazhko behaviour was found in 18 stars, which were supposed to be modulated. However, five of these eighteen stars are certain Blazhko variables, but with too small amplitude of modulation to be detectable in ASAS or SW data.  The incidence rate of RRab Blazhko stars from our analysis is at least 31~\\%, which is a much higher percentage than given by previous studies based on data from surveys with similar data quality. This increase would be logically expected, since more time-extended datasets were examined (in many cases simultaneously in ASAS and SW surveys). However, it is still less than the incidence rates determined in the framework of precise ground-base and space surveys. With certainty there remained undetected low-amplitude Blazhko variables, as well as stars with the shortest and the longest modulation periods, simply due to nature of the data from ASAS and SW surveys.  A surprising result was a very high incidence of stars with some peculiarity in their Blazhko effect -- each eighth star of our Blazhko sample showed signs of the changes in Blazhko effect or compound modulation. An interesting examples are IK Hya and RZ CVn. Both stars probably undergo changes in their pulsation and modulation characteristics. IK Hya could be another star with a few-year-long cycle analogical to four-year cycle in RR~Lyrae itself. Five of ten identified multiple-modulated stars showed the ratio between their modulation periods in ratio of small integers.  Our study confirmed that classification and statistics of modulated stars according to their side peaks, which are based on low-quality data, are not reliable. It was also shown that the sensitivity of automatic procedures used to identify the Blazhko effect is insufficient -- one fourth of the Blazhko variables were identified for the first time. It is likely that the incidences and characteristics of Blazhko variables identified in other surveys \\citep[e.g. in LMC based on OGLE, ][]{alcock2003} could be slightly different than proposed. Our conclusions should serve as a challenge for developing better tools for automatic data analysis, because a manual examination of\tmany thousand stars is impossible to perform."
    },
    "1402/1402.1329_arXiv.txt": {
        "abstract": "We study an isothermal system of semi-degenerate self-gravitating fermions in general relativity. Such systems present mass density solutions with a central degenerate core, a plateau and a tail which follows a power law behaviour $r^{-2}$. The different solutions are governed by the free parameters of the model: the degeneracy and temperature parameters at the center, and the particle mass $m$. We then analyze in detail the free parameter space for a fixed $m$ in the keV regime, by studying the one-parameter sequences of equilibrium configurations up to the critical point, which is represented by the maximum in a central density ($\\rho_0$) Vs. core mass ($M_c$) diagram. We show that for fully degenerate cores, the known expression for the critical core mass $M_c^{cr}\\propto m_{pl}^3/m^2$ is obtained, while instead for low degenerate cores, the critical core mass increases showing the temperature effects in a non linear way. The main result of this work is that when applying this theory to model the distribution of dark matter in galaxies from the very center up to the outer halos, we do not find any critical core-halo configuration of self-gravitating fermions, which be able to explain the super massive dark object in their centers together with an outer halo simultaneously. ",
        "introduction": "Systems of self-gravitating semi-degenerate fermions in general relativity were studied in \\cite{rr} and more recently with applications to dark matter in galaxies in \\cite{charly}. It was shown that, for a given central temperature parameter ($\\beta_0$) in agreement with the corresponding observed halo circular velocity, there are lower bounds for the central degeneracy parameter ($\\theta_0$) and particle mass ($m\\gtrsim0.4$ keV) above which the \\textit{observed} halo mass and radius are fulfilled. The density profiles solutions in this approach present a novel core-halo morphology composed by a quantum degenerate core followed by a low degenerate plateau until they reach the $r^{-2}$ Boltzmannian regime. This interesting overall morphology provides the flat rotation curves in the outermost part of the galaxies as well as a possible alternative to massive black holes in their centers (see \\cite{rrchina} and \\cite{firstwork}).  The system of Einstein equations are written in a spherically symmetric space-time metric $g_{\\mu \\nu}={\\rm diag}(e^{\\nu},-e^{\\lambda},-r^2,-r^2\\sin^2\\theta)$, where $\\nu$ and $\\lambda$ depend only on the radial coordinate $r$, together with the thermodynamic equilibrium conditions of Tolman \\cite{tolman}, and Klein \\cite{klein}, \\begin{equation} e^{\\nu/2} T=const.\\, , \\quad e^{\\nu/2}(\\mu+m c^2)=const, \\nonumber \\end{equation} where $T$ is the temperature, $\\mu$ the chemical potential, $m$ the particle mass and $c$ the speed of light. We then write the system of Einstein equations in the following dimensionless way, \\begin{align} &\\frac{d\\hat M}{d\\hat r}=4\\pi\\hat r^2\\hat\\rho \\label{eq:1}\\\\ &\\frac{d\\theta}{d\\hat r}=\\frac{\\beta_0(\\theta-\\theta_0)-1}{\\beta_0} \\frac{\\hat M+4\\pi\\hat P\\hat r^3}{\\hat r^2(1-2\\hat M/\\hat r)}\\\\ &\\frac{d\\nu}{d\\hat r}=\\frac{\\hat M+4\\pi\\hat P\\hat r^3}{\\hat r^2(1-2\\hat M/\\hat r)} \\\\ &\\beta_0=\\beta(r) e^{\\frac{\\nu(r)-\\nu_0}{2}}\\, . \\label{eq:2} \\end{align} The variables of the system are the mass $M$, the metric factor $\\nu$ , the temperature parameter $\\beta=k T/(m c^2)$ and the degeneracy parameter $\\theta=\\mu/(k T)$. The dimensionless quantities are: $\\hat r=r/\\chi$, $\\hat M=G M/(c^2\\chi)$, $\\hat\\rho=G \\chi^2\\rho/c^2$ and $\\hat P=G \\chi^2 P/c^4$, with $\\chi=2\\pi^{3/2}(\\hbar/mc)(m_p/m)$ and $m_p=\\sqrt{\\hbar c/G}$ the Planck mass.  The mass density $\\rho$ and pressure $P$ are given by Fermi-Dirac statistics (see also \\cite{charly}).  This system is solved for a fixed particle mass $m$ in the keV range, with initial conditions $M(0)=\\nu(0)=0$, and given parameters $\\theta_0>0$ (depending on the chosen central degeneracy), and $\\beta_0$. We thus construct a sequence of different thermodynamic equilibrium configurations where each point in the sequence has different central temperatures $T_0$ and central chemical potential $\\mu_0$, so that satisfy the $\\theta_0$ fixed condition.  Defining the core radius $r_c$ of each equilibrium system at the first maximum of its rotation curve, or equivalently at the degeneracy transition point in which $\\theta(r_c)=0$, we represent the results obtained for each sequence in a central density ($\\rho_0$) vs. core mass ($M_c$) diagram (see Fig.\\ref{fig:1}). It is shown that the critical core mass $M_c^{cr}$ is reached at the maximum of each $M_c(\\rho_0)$ curve.  \\begin{figure}[!hbtp] \\centering \\includegraphics[width=\\linewidth]{criticalMassT.eps} \\caption{Different sequences of equilibrium configurations plotted in a central density ($\\rho_0$) Vs. core mass ($M_c$) diagram. The critical core mass is reached at the maximal value of $M_c$. Each sequence is built for selected values of $\\theta_0=\\mu_0/kT_0$ and different values of $T_0, \\mu_0$ varying accordingly.} \\label{fig:1} \\end{figure}  It is important to emphasize that we are not interested in follow the history of equilibrium states of one specific system. Thus, the standard stability analysis as done for compact stars or in dense stellar cluster (see e.g. \\cite{rrcluster}), which is based on the constancy of the entropy per nucleon (S/N) along the equilibrium sequence of a given configuration, does not apply here. Nonetheless, in computing the $M_c(\\rho_0)$ curves in Fig.~\\ref{fig:1} we have explored the full range of $\\theta_0>1$ and $\\beta_0>10^{-10}$ parameters (including the critical ones). Then the equilibrium sequences with constant specific entropy (S/N), which differ from the ones with constant $\\theta_0$ considered here, necessarily must be contained within the full ($T_0,\\mu_0$) parameter space covered in Fig.~\\ref{fig:1}.  In Table~\\ref{table:1} we show a set of central critical parameters of the model together with the correspondent critical core masses, for a very wide range of fixed central degeneracy parameters $\\theta_0$ and $m=8.5$ keV$/c^2$.  \\begin{table} \\begin{ruledtabular} \\begin{tabular}{cccc} $\\theta_0$ & $\\beta_0^{cr}$ & $\\mu_0^{cr}/mc^2$ & $M_c^{cr} (M_\\odot)$ \\\\ \\colrule 1 & $6.45\\times10^{-2}$ & $6.45\\times10^{-2}$ & $1.59\\times10^{10}$ \\\\ 5 & $2.23\\times10^{-2}$ & $1.11\\times10^{-1}$ & $7.91\\times10^9$ \\\\ 40 & $8.33\\times10^{-3}$ & $3.33\\times10^{-1}$ & $7.44\\times10^9$ \\\\ 55 & $6.06\\times10^{-3}$ & $3.33\\times10^{-1}$ & $7.44\\times10^9$ \\\\ 100 & $3.33\\times10^{-3}$ & $3.33\\times10^{-1}$ & $7.44\\times10^9$ \\\\ \\end{tabular} \\caption{Critical temperature parameter and normalized chemical potential at the center of each different critical configuration, for different fixed central degeneracies.} \\end{ruledtabular} \\label{table:1} \\end{table}  Defining the halo radius of each configuration at the onset of the flat rotation curve, we show in Table II the critical halo magnitudes $r_h^{cr}$, $M_h^{cr}$ and $v_h^{cr}$ corresponding to the same set of critical parameters as given in Table I.  \\begin{table} \\begin{ruledtabular} \\begin{tabular}{cccc} $\\theta_0$ & $r_h^{cr} (pc)$ & $M_h^{cr}/mc^2 (M_\\odot)$ & $v_h^{cr} (km/s)$ \\\\ \\colrule 1 & $4.4\\times10^{-1}$ & $5.7\\times10^{11}$ & $7.5\\times10^4$ \\\\ 5 & $4.0\\times10^{-1}$ & $4.3\\times10^{11}$ & $6.2\\times10^4$ \\\\ 40 & $4.3\\times10^{3}$ & $1.1\\times10^{15}$ & $3.3\\times10^4$ \\\\ 55 & $2.9\\times10^{5}$ & $6.0\\times10^{16}$ & $2.9\\times10^4$ \\\\ 100 & $2.0\\times10^{11}$ & $2.3\\times10^{22}$ & $2.2\\times10^4$ \\\\ \\end{tabular} \\caption{Critical halo magnitudes of different critical configurations, for different fixed central degeneracies as given in Table I.} \\end{ruledtabular} \\label{table:2} \\end{table}  The results obtained in Tables I and II imply a marked division in two different families depending on the value of $M_c^{cr}$.  \\textit{i}) The first family: the critical mass has roughly a constant value $M_c^{cr}=7.44\\times10^9 M_\\odot$. This family corresponds to large values of the central degeneracy ($\\theta_0\\geq40$), where the critical temperature parameter is always lower than $\\beta_0^{cr}\\lesssim8\\times10^{-3}$ and the critical chemical potential $\\mu_0^{cr}\\approx$ const. Physically, these highly degenerate cores are entirely supported against gravitational collapse by the degeneracy pressure. In this case the critical core mass is uniquely determined by the particle mass according the relation $M_c^{cr}\\propto m_{pl}^3/m^2$ (see also section III).  \\textit{ii}) The second family: the critical core mass increases from $M_c^{cr}=7.44\\times10^9 M_\\odot$ up to $M_c^{cr}\\sim10^{10} M_\\odot$. This case corresponds to critical cores with a lower central degeneracy compared with the former family ($1<\\theta_0<40$). Here the critical temperature parameter ($\\beta_0\\sim10^{-2}$), is closer to the relativistic regime with respect to the first family. This result physically indicates that the thermal pressure term has now an appreciable contribution to the total pressure, which supports the critical core against gravitational collapse. In this case $M_c^{cr}$ is completely determined by the particle mass $m$, the central temperature $T_0^{cr}$ and the central chemical potential $\\mu_0^{cr}$ (see section III).  In Figs.~(\\ref{fig:2}) and (\\ref{fig:3}) we show a critical metric factor $e^{\\nu/2}$ and a critical temperature $kT$ as function of the radius for the two different families mentioned above.  \\begin{figure} \\centering \\includegraphics[width=.9\\linewidth]{nutemphot.eps} \\caption{The critical temperature profile of the system (in keV) and the critical metric, for $\\theta_0=5$ and $\\beta_0^{cr}=2.23\\times 10^{-2}$. The dashed line corresponds to the isothermality condition, $Te^{\\nu/2}=const$.} \\label{fig:2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=.9\\linewidth]{nutemp.eps} \\caption{The critical temperature of the system (in keV) and the critical metric, for $\\theta_0=55$ and $\\beta_0^{cr}=6.06\\times10^{-3}$. The red line corresponds to the isothermality condition, $Te^{\\nu/2}=const$.} \\label{fig:3} \\end{figure} ",
        "conclusions": ""
    },
    "1402/1402.5166_arXiv.txt": {
        "abstract": "Inferences of sub-surface flow velocities using local domain ring-diagram helioseismology depend on measuring the frequency splittings of oscillation modes seen in acoustic power spectra. Current methods for making these measurements utilize maximum-likelihood fitting techniques to match a model of modal power to the spectra. The model typically describes a single oscillation mode, and each mode in a given power spectrum is fit independently. We present a new method that produces measurements with greater reliability and accuracy by fitting multiple modes simultaneously. We demonstrate how this method permits measurements of sub-surface flows deeper into the Sun while providing higher uniformity in data coverage and velocity response closer to the limb of the solar disk. While the previous fitting method performs better for some measurements of low-phase-speed modes, we find this new method to be particularly useful for high phase-speed modes and small spatial areas. ",
        "introduction": "\\label{sec:intro}  Helioseismology determines the structure and dynamics of the solar interior through analysis of seismic waves observed at the surface. Ring-diagram helioseismology (\\opencite{hill_1988}; \\opencite{basu_1999}; \\opencite{haber_2002}) investigates sub-surface horizontal flows by measuring the direction-dependent frequency shift of oscillation modes. The Doppler shift of the frequency of an oscillation mode due to a sub-surface flow is expressed as \\begin{equation} \\delta \\omega_{n}(\\bvec{k}) = \\bvec{k} \\cdot \\bvec{u}_{n}(k), \\label{eqn:freqshift} \\end{equation} where $\\bvec{k}$ and $n$ are the horizontal wavenumber and radial order of the oscillation mode, and $\\bvec{u}_{n}(k)$ is a spatial average of the horizontal velocity within the Sun, \\begin{equation} \\bvec{u}_{n}(k) = \\langle \\bvec{v}(\\bvec{r}) \\rangle = \\int \\! K_{n}(\\bvec{r};k) \\ \\bvec{v}(\\bvec{r}) \\ \\mathrm{d}^{3}r. \\label{eqn:kerneldef} \\end{equation} Here, $\\bvec{v}(\\bvec{r})$ is the true horizontal sub-surface flow velocity at any point \\noparen{$\\bvec{r}$} in the Sun, and $K_{n}(\\bvec{r};k)$ is the weighting function\\emdash or sensitivity kernel\\emdash associated with each mode ($k$,$n$), which describes the spatial extent over which the true velocity is averaged to create a single frequency shift \\cite{birch_2007}. The variation of the frequency shift as a function of direction (the frequency splitting) is measured to provide an estimate of $\\bvec{u}_{n}(k)$.  While the interpretation of the frequency splitting is straightforward, the method of extracting it from the data is not. Traditionally, a model of the spectral power is fit to oscillation modes visible in power spectra of line-of-sight velocity observed in the photosphere. The model accounts for a frequency shift of the modal power as a function of horizontal direction, and this shift is directly translated into a velocity as shown in Equation \\eqnref{eqn:freqshift}. The specifics of the fitting procedure determine how well the frequency splittings are measured, as well as what other qualities of the power spectra are taken into account. There are currently two commonly used fitting procedures in the \\textit{Helioseismic and Magnetic Imager} (HMI) Ring-Diagram Pipeline \\cite{bogart_2011a}. The first method considered in this article is one introduced by \\inlinecite{haber_2002}, which fits a frequency-shifted Lorentzian model to individual modes. Since this method analyzes single ridges of modal power sequentially, we refer to it as the Single-Ridge Fitting method (SRF). The second method used in the HMI Pipeline, which will not be considered here, also fits modes independently \\cite{basu_1999}, but uses a significantly different power model that includes asymmetries in modal power, two background terms, and a host of other wavenumber- and direction-dependent modifications.  In this article, we present a new fitting method that utilizes a model similar to that used in the SRF method, but modified to permit multiple radial orders to be fit simultaneously. The development of this Multi-Ridge Fitting (MRF) method is an attempt to improve upon the performance of the SRF method in terms of reliability and accuracy of measured frequency splittings. We present a comparison of the frequency splittings (interpreted as an average velocity) from a common data set processed with both fitting methods. We focus on the performance for measurements of modes that reach deepest into the Sun and for measurements made near the solar limb. As metrics of reliability and accuracy we consider the fit success rate, the typical random errors in the measurements, and how well each method recovers a known velocity.  In Section \\ref{sec:data} we outline the data-processing steps taken to prepare solar observations for helioseismic mode fitting. In Section \\ref{sec:fitting} we describe the SRF and MRF fitting methods and how they differ in procedure. In Section \\ref{sec:comparison} we compare the performance of the two fitting methods using a common data set. In Section \\ref{sec:discussion} we discuss the implications of the results in the context of improving accuracy and data coverage in ring-diagram helioseismology. ",
        "conclusions": "\\label{sec:discussion}  Extracting frequency shifts that are interpreted as sub-surface velocities is done through fitting a model of acoustic power to a three-dimensional power spectrum. By fitting a model that includes various effects seen in the data to different segments of the three-dimensional spectrum, many unique fitting methods can be created. We have presented two such fitting procedures that differ both in the model used and the selection of data considered in a single optimization. These methods use identical expressions for the velocity-induced frequency splitting for each oscillation mode, yet produce significantly different results for the velocity measurements themselves.  \\subsection{Improvements in Depth for Inversions}  The MRF method is able to obtain a greater number of velocity measurements for higher phase-speed modes in a given tile size than the SRF method. The increase in the success rate is also prominent for small tile sizes, suggesting that the treatment of mode blending plays an important role. Smaller tiles exhibit more leakage of modal power across wavenumbers than larger tiles due to spatial apodization. The spreading of power increases the effective width of modes at a single wavenumber. High-phase-speed modes are already spaced closer in wavenumber than other modes, enhancing the mode blending for their region of each spectrum. The SRF model does not accommodate any overlap of power from neighboring modes when fitting a single mode, causing a disparity between model and data that worsens at higher phase-speeds and smaller tile sizes. The MRF method, in accounting for mode overlap, returns a higher number of successful measurements for these cases that can be used for studying deep flows and high horizontal resolution flows.  Referring to Equation \\eqnref{eqn:kerneldef}, the set of measurements obtained from a single tile \\noparen{$\\bvec{u}_n(k)$} are measures of the true sub-surface velocity \\noparen{$\\bvec{v}(\\bvec{r})$} within a single region \\noparen{$d^3r$} weighted with different sensitivity kernels \\noparen{$K(\\bvec{r})$}. The sensitivity kernels generally do not have sufficient isolation of sensitivity at any given depth, so while having a larger number of high-phase-speed measurements results in reaching deeper into the Sun, these measurements alone do not provide a localized determination of sub-surface flows. The velocity measurements from many modes can be merged through an inversion to produce an estimate of the true sub-surface flow velocity in an isolated region. After finding a linear combination of sensitivity kernels that produces a single isolated peak of sensitivity at a target location, the estimated velocity for that region is constructed using the same linear combination of the associated velocity measurements: \\begin{equation} \\bvec{w} = \\sum_{i} a_{i} \\bvec{u}_{i}. \\label{eqn:invsoln} \\end{equation} \\noindent Here $\\bvec{w}$ is the estimated sub-surface flow velocity at a selected target point, $\\bvec{u_{i}}$ are the velocity measurements obtained from each mode $(n,k)$, and $a_{i}$ are the coefficients of the linear combination determined by the inversion. To estimate the localization of each solution point, we can construct an averaging kernel that represents the spatial distribution of sensitivity for a solution point, just as a sensitivity kernel does for measured velocities: \\begin{equation} H(\\bvec{r}) = \\sum_{i} a_{i} K_{i}(\\bvec{r}). \\label{eqn:avgkercombination} \\end{equation} The set of modes that we are able to fit reliably using either method determine the basis set of kernels that can be used to construct a locally isolated averaging kernel $H(\\bvec{r})$. It is not only how high in phase-speed the fitting method can reach, but also how reliably it can obtain successful measurements at all phase-speeds, that determines how well we can constrain sub-surface flows.  To determine how well averaging kernels can be constructed using the mode set available from each fitting method operating on $16^{\\circ}$ tiles, we use a simple 1D optimally localized averaging (OLA) inversion (\\opencite{dalsgaard_1990}; \\opencite{schou_1998}). The target functions are Gaussian profiles in depth and we restrict the set of sensitivity kernels available for inversion to those with a $\\mathrm{80\\,\\%}$ or higher success rate. Figure \\ref{fig:kernels} compares averaging kernels constructed from each kernel set to the target functions used. \\begin{figure}[] \\centering \\includegraphics{kernels2_16.ps} \\caption{(a) Target functions for the 1D OLA inversion. (b) Averaging kernels constructed from sensitivity kernels from the SRF mode set. (c) Averaging kernels constructed from sensitivity kernels from the MRF mode set. The colors indicate different target depths. The SRF method is only able to produce averaging kernels down to a depth of about $30 \\ \\mathrm{Mm}$, while the MRF method reaches down to $45 \\ \\mathrm{Mm}$.} \\label{fig:kernels} \\end{figure}  For shallow targets, both sets are able to accurately recreate the Gaussian target function with minimal sidelobes in depth. As the target depth increases, the number of kernels available for constructing averaging kernels decreases. Past a target depth of 30 Mm, the SRF set does not contain enough kernels to create a sufficiently isolated peak of sensitivity. The MRF set still contains enough unique kernels that reach past this depth to construct reasonable averaging kernels down to 45 Mm.   While the ability to construct reasonable averaging kernels through inversion is crucial for interpreting sub-surface flow results, this ignores the influence of measured errors. For OLA inversions, an optimization is done to balance the shape of the averaging kernel against the propagated error: \\begin{equation} \\sigma^2 = \\sum_{i} ( a_{i} \\sigma_{i} )^2. \\label{eqn:inverr} \\end{equation} In this way, it is both the available mode set and the associated errors that determine the usefulness of the inversion results. Figure \\ref{fig:inverr} shows the propagated inversion error as a function of the averaging kernel center-of-gravity depth for each fitting method. To mimic the result of inverting an average of many days of data over all longitudes to look at large-scale mean flows, the error for each mode has been scaled: \\begin{equation} \\sigma_i \\rightarrow \\frac{\\sigma_{i}}{\\sqrt{N_d N_l r_i}}, \\label{eqn:inverr2} \\end{equation} where $r_i$ is the success rate for a given mode, $N_d$ is the number of days to average over, and $N_l$ is the number of longitudes to average over. We set $N_d$ to 180 to mimic inversions of six months of data and $N_l$ to 17 for the typical number of tiles spanning the solar disk in longitude. The various curves for different regions of the solar disk demonstrate how much the dependence of the success rate and average error on disk position affect inversion results. The sharp downturn seen at the deep end of each curve indicates where the averaging kernels no longer resemble the target functions and are therefore useless.  Due to the higher success rate for each mode, the MRF method is able to produce useful inversion results to a greater depth than the SRF method. As the distance from disk center increases, the MRF results remain mostly constant while the SRF results start to deteriorate both in error magnitude and kernel depth. Despite larger average error in the MRF results for shallow modes seen in Figure \\ref{fig:avgerr}, the higher success rate for these modes causes the inversion errors to not show a significant increase relative to the SRF results.  \\begin{figure}[] \\centering \\includegraphics{inverr.ps} \\caption{Typical error magnitudes propagated through 1D OLA inversion as a function of depth for multiple distances from disk center. The sharp turnoff at the deep end of each curve demonstrates where proper averaging kernels can no longer be produced. The inversions done with SRF results tend to turn off between 33 and 35 Mm depth while those done with MRF results turn off between 38 and 40 Mm.} \\label{fig:inverr} \\end{figure}   \\subsection{Spatial Variability and Implications}  Anisotropy along the azimuthal direction of the power spectrum appears to be the primary cause of fitting problems near the solar limb. The SRF method tackles this issue by operating on spectra that have been ``flattened\" along azimuth. This step not only amplifies the modal power in certain directions in order to eliminate anisotropy, but also amplifies the noise. By accounting for the natural variation of power along each ridge as well as in the background power, the MRF method is able to push closer to the solar limb. As seen in Figure \\ref{fig:successdisk}, the mode set measureable by the SRF method is position-dependent. The higher-order modes see a drop in success rate closer to disk center than lower-order modes. The depth to which one can measure is then shallower near the limb. The MRF method, in contrast, allows for consistent determinations of sub-surface flows within $75^{\\circ}$ of disk center.  Large tile sizes tend to provide the most accurate velocity measurements for both fitting methods, with a slight tendency for both to overestimate at high latitudes (Figure \\ref{fig:track4}). Both methods exhibit spatial variability in the velocity response for smaller tile sizes, but only for low-phase-speed modes. The MRF response is radially symmetric with the worst velocity response occurring near disk center, while the worst regions of the SRF response appear far from disk center and have opposite signs between the north--south and east--west direction. While the magnitude of this variability is relatively small, there are significant implications when attempting to analyze small residual flows left over from the subtraction of large-scale flows. Tiles at high latitudes tracked at the Carrington rate see an overall flow speed of nearly $200 \\ \\mathrm{m \\ s^{-1}}$ due to differential rotation. The systematic underestimation seen in the SRF method introduces an anomalous retrograde flow of around $20 \\ \\mathrm{m \\ s^{-1}}$, similar to the magnitude of the center-to-limb velocity systematic for low-phase-speed modes \\cite{greer_2013}. It is unclear what causes the systematic inaccuracy in either fitting method, although the strong dependence on disk position suggests a possible coupling between velocity measurements and power anisotropy.  As mentioned in Section \\ref{sec:mrf}, there is a significant increase in the computational cost when switching from the SRF method to the MRF method. While processing a day's worth of tiles with the SRF method is typically one of the fastest steps in the HMI Pipeline, the MRF method brings the process of fitting more in line with the time needed for tracking. It is important to consider when this additional computational burden is justified. The SRF method produces high-quality results for low-phase-speed modes and for large tiles near disk center. This leaves three distinct cases where the MRF method provides improvements: measuring higher-phase-speed modes, pushing closer to the limb, and using small tiles. The implications of this new procedure on determining sub-surface flows follow these three technical improvements. For a given tile size, we are able to extend our analysis deeper into the Sun while maintaining a constant horizontal and temporal resolution. Analysis to these extended depths can be performed consistently across most of the solar disk, providing uniform coverage over a larger fraction of the Sun. The increased reliability of small tile sizes permits higher horizontal resolution analysis of sub-surface flows.  \\textbf{Acknowledgements} This work was supported by NASA through NASA grants NNX08AJ08G, NNX08AQ28G, and NNX09AB04G. The data used here are courtesy of NASA/SDO and the HMI science team. SDO is a NASA mission, and the HMI project is supported by NASA contract NAS5-02139."
    },
    "1402/1402.5399_arXiv.txt": {
        "abstract": "Very light WIMPs ($\\chi$), thermal relics that annihilate late in the early Universe, change the energy and entropy densities at BBN and at  recombination.  BBN, in combination with the CMB, can remove some of the degeneracies among light WIMPs and equivalent neutrinos, constraining the existence and properties of each.  Depending on the nature of the light WIMP (Majorana or Dirac fermion, real or complex scalar) the joint BBN + CMB analyses set {\\bf lower} bounds to \\mchi~in the range $0.5 - 5\\,{\\rm MeV}$ ($m_{\\chi}/m_{e} \\gsim 1 - 10$), and they identify {\\bf best fit} values for \\mchi~in the range $5 - 10\\,{\\rm MeV}$.  The {\\bf joint} BBN + CMB analysis finds a {\\bf best fit} value for the number of equivalent neutrinos, \\Deln~$\\approx 0.65$, nearly independent of the nature of the WIMP.   In the absence of a light WIMP (\\mchi~$\\gsim 20\\,{\\rm MeV}$), \\neff~$= 3.05(1 + \\Deln/3)$.  In this case, there is excellent agreement between BBN and the CMB, but the joint fit reveals \\Deln~$= 0.40\\pm0.17$, disfavoring standard big bang nucleosynthesis (SBBN) (\\Deln~= 0) at $\\sim 2.4\\,\\sigma$, as well as a sterile neutrino (\\Deln~= 1) at $\\sim 3.5\\,\\sigma$.   The best BBN +  CMB joint fit disfavors the absence of dark radiation (\\Deln~= 0 at $\\sim 95\\%$ confidence), while allowing for the presence of a sterile neutrino (\\Deln~= 1 at $\\lsim 1\\,\\sigma$).  For all cases considered here, the lithium problem persists.  These results, presented at the TAUP 2013 Conference, are based on \\citet{kngs}. ",
        "introduction": "\\label{intro}  In the absence of ``extra\", equivalent neutrinos (dark radiation) or light ($\\lsim 20\\,{\\rm MeV}$), weakly interacting massive particles (WIMPs), the particle content relatively late in the early Universe is quite simple.  After the \\epm pairs (and all the other more massive standard model (SM) particles) have annihilated ($T \\lsim m_{e}$), the only remaining SM particles are the CMB photons and the three SM neutrinos ($\\nu_{e}$, $\\nu_{\\mu}$, $\\nu_{\\tau}$).  At these early epochs the Universe is ``radiation dominated\" and the energy density may be written as $\\rho_{\\rm R} = \\rho_{\\gamma} + 3\\,\\rho_{\\nu}$, where $3\\,\\rho_{\\nu}$ accounts for the contributions from the three SM neutrinos.  During these early epochs the contributions to the total mass/energy density from the baryons (B) and the dark matter (DM), as well as any dark energy (DE), are very subdominant compared to $\\rho_{\\rm R}$.  More generally, in addition to the SM neutrinos, there may be extra, ``beyond the standard model\" particles that, like the SM neutrinos, are extremely light ($\\lsim 10\\,{\\rm eV}$) and very weakly interacting.  During the early (or, even, relatively late) evolution of the Universe these neutrino-like particles, so called ``equivalent neutrinos\", will contribute to the energy density.  The energy density controls the early Universe expansion rate.  If \\Deln~counts the contribution of equivalent neutrinos, often referred to as ``dark radiation\", $\\rho_{\\rm R} = \\rho_{\\gamma} + (3 + \\Deln)\\,\\rho_{\\nu}$.  The contribution to \\Deln~of an equivalent neutrino that decouples along with the SM neutrinos (at $T = T_{\\nu d}$) will be \\Deln~= 1 for a Majorana fermion (\\eg, a sterile neutrino), \\Deln~= 2 for a Dirac fermion or, \\Deln~= 4/7 for a real scalar.  In general, \\Deln~is an integer (fermions) or an integer multiple of 4/7 (bosons).  However, an equivalent neutrino that is more weakly interacting than the SM neutrinos, will have decoupled earlier in the evolution of the Universe and its contribution to \\Deln~will be suppressed by the heating of the SM neutrinos (and photons) when the heavier SM particles decay and/or annihilate.  Therefore, in principle, there is no reason that \\Deln~should be an integer or an integer multiple of 4/7 (for further discussion see \\citet{steig13}; for a specific example of three, very weakly coupled, right-handed neutrinos, see \\citet{haim} and for the example of a weakly coupled scalar particle see \\citet{weinberg}).  After the SM neutrinos decouple, when $T = T_{\\nu d} \\approx 2 - 3\\,{\\rm MeV}$, the \\epm pairs annihilate, heating the photons but not the neutrinos.  Prior to neutrino decoupling (and \\epm annihilation), the neutrinos, \\epm pairs, and the photons are in equilibrium at the same temperature, $T_{\\nu} = T_{e} = T_{\\gamma}$ but, after \\epm annihilation, the photons are hotter than the relic neutrinos (or, equivalently, the neutrinos are cooler than the photons).  In most simplified analyses it is assumed that the neutrinos decoupled instantaneously and that the electrons were effectively massless at neutrino decoupling.  With these approximations the late time (after \\epm annihilation is complete) ratio of neutrino and photon temperatures, $(T_{\\nu}/T_{\\gamma})_{0} = (4/11)^{1/3}$, follows from entropy conservation.  The late time ratio of the energy density in one species of neutrino ($\\rho^{0}_{\\nu}$) to that in the photons is $(\\rho^{0}_{\\nu}/\\rho_{\\gamma})_{0} = 7/8\\,(T_{\\nu}/T_{\\gamma})^{4}_{0} = 7/8\\,(4/11)^{4/3}$.  However, at neutrino decoupling $m_{e}/T_{\\nu d} \\approx 0.2 \\neq 0$ and, as a result, $\\rho_{\\nu}$ differs (by a small amount) from $\\rho^{0}_{\\nu}$ \\citep{steig13}.  Furthermore, the neutrinos don't really  decouple instantaneously.  While the neutrinos are partially coupled to the annihilating \\epm pairs they share a small amount of the energy released by the annihilation \\citep{dolgov,dolgov-osc,enqvist,hannestad,mangano}.  These effects can be accounted for by introducing \\neff, the ``effective number of neutrinos\" where, at late times ($T_{0} \\ll m_{e}$), $\\rho_{\\rm R\\,0} \\equiv \\rho_{\\gamma\\,0} + {\\rm N}_{\\rm eff}\\,\\rho^{0}_{\\nu\\,0}\\,,$ \\beq {\\rm N}_{\\rm eff} = 3\\,\\bigg[{11 \\over 4}\\bigg({T_{\\nu} \\over T_{\\gamma}}\\bigg)^{3}_{0}\\bigg]^{4/3}\\bigg(1 + {\\Deln \\over3}\\bigg)\\,. \\eeq It should be kept in mind that while the relative contributions of neutrinos and photons to the total radiation density may be evolving before and during BBN, \\neff~is a ``late time\" quantity, evaluated long after BBN has ended, when the only relativistic particles remaining are the photons and the neutrinos.  Under the assumptions of instantaneous neutrino decoupling and $m_{e} \\ll T_{\\nu d}$, \\neff~= 3 + \\Deln.  Keeping the instantaneous decoupling approximation but correcting for the finite electron mass, \\neff~$\\approx 3.02(1 + \\Deln/3)$ \\citep{steig13}.  Following the non-instantaneous neutrino decoupling and allowing for the finite electron mass, \\neff~$\\approx 3.05(1 + \\Deln/3)$ \\citep{mangano}.  It should be noted that in this latter case there is, in addition, a very small, but not entirely negligible correction to the BBN predicted primordial helium abundance \\citep{mangano}.  Since the expansion rate (Hubble parameter) of the radiation dominated early Universe is $H \\propto \\rho_{\\rm R}^{1/2}$, the presence of dark radiation (\\Deln~$\\geq 0$) results in a speed up to the early Universe expansion rate.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=0.45\\columnwidth]{neff0_vs_mchi.pdf} \\hskip .4in \\includegraphics[width=0.45\\columnwidth]{neff_vs_mchi.pdf} \\caption{The left panel shows N$^{0}_{\\rm eff}$ (\\neff~when \\Deln~= 0) as a function of the WIMP mass for electromagnetically coupled light WIMPs in the absence of equivalent neutrinos.  From bottom to top, the solid red curve is for a Dirac WIMP,  the dashed green curve is for a complex scalar, the solid black curve is for a Majorana fermion and, the dashed blue curve is for a real scalar.  The horizontal, red/pink bands are the Planck CMB 68\\% and 95\\% allowed ranges for N$_{\\rm eff}$.  The right panel specializes to the case of a Majorana fermion WIMP, showing N$_{\\rm eff}$ as a function of the WIMP mass for \\Deln~equivalent neutrinos.  The solid  curve is for $\\Delta {\\rm N}_{\\nu} = 0$, the short dashed curve is for $\\Delta {\\rm N}_{\\nu} = 1$ and, the long dashed curve is for $\\Delta {\\rm N}_{\\nu} = 2$.  The horizontal red bands are the Planck CMB 68\\% and 95\\% allowed ranges for N$_{\\rm eff}$, including baryon acoustic oscillations in the CMB constraint.} \\label{fig:neffvsmchi} \\end{center} \\end{figure}  So far, the possibility of a very light, weakly interacting, massive particle, a WIMP $\\chi$, has been ignored in the discussion here.  The difference between a WIMP and an equivalent neutrino is that a WIMP remains thermally coupled to the SM particles until after it has become non-relativistic and begins annihilating.  As a result, the light WIMP annihilation heats the remaining SM particles (either the photons and, possibly, the \\epm pairs if the WIMP couples electromagnetically or, the SM neutrinos if the WIMP only couples to them).  Note that in the analysis and discussion here, the WIMP {\\bf need not} be a dark matter candidate; the WIMP could be a sub-dominant component of the dark matter ($\\Omega_{\\chi} < \\Omega_{\\rm CDM}$).  For example, the WIMP could be a light, millicharged particle such as that proposed by \\citep{millicharged} and discussed by Dolgov at this conference.  Here, as in \\citet{kngs}, we specialize to the case of a light WIMP coupled only to the photons and \\epm pairs.  The relevant role for BBN and the CMB played by such a light WIMP is that its annihilation heats the photons relative to the decoupled SM neutrinos, changing (reducing) $(T_{\\nu}/T_{\\gamma})_{0}$.  After \\epm and WIMP annihilation, at fixed photon temperature, the neutrinos, SM and equivalent, are cooler than the photons.  In this case, \\neff~is a function of \\mchi~(see \\citet{steig13} and references therein).  Since the expansion rate of the early Universe is controlled by the energy density, any modification of \\neff~will be reflected in a non-standard expansion rate (\\eg, during BBN and at recombination).  In addition, extremely light WIMPs (\\mchi~$\\lsim m_{e}$) will annihilate so late that, if their annihilation produces photons, they will modify the baryon-to-photon ratio ($\\eta_{10} = 10^{10}(n_{\\rm B}/n_{\\gamma})_{0} = 273.9\\,\\omb$) during or after BBN.  BBN can probe \\neff~(through the effects of the neutrinos on the expansion rate and on the weak rates that regulate the neutron to proton ratio) as well as the universal ratio of baryons-to-photons.  At late times in the early Universe, \\eg, at recombination, the CMB can also probe \\omb~and \\neff.  As independent probes of the effective number of neutrinos (\\neff) or the number of equivalent neutrinos (\\Deln) and the universal baryon density (\\omb~or $\\eta_{10}$), BBN and the CMB can help to break the degeneracies among these parameters and the WIMP mass (and spin/statistics) and to constrain their allowed ranges (see, \\citet{steig13} \\& \\citet{kngs} and Fig.\\,\\ref{fig:neffvsmchi}).  \\subsection{Planck CMB Constraints}  In their analysis, \\citet{kngs}, whose results are described and summarized here, adopted the CMB constraints on \\omb~and \\neff~from the Planck $\\Lambda\\mathrm{CDM}+N_\\mathrm{eff}$ fit including supplementary baryon acoustic oscillation (BAO) data \\citep{planck}.  The correlations between these quantities were included in \\citep{kngs} and in the analysis here.  From the CMB we adopted \\omb~$= 0.0223 \\pm 0.0003$ ($\\eta_{10} = 6.11 \\pm 0.08$) and \\neff~$= 3.30 \\pm 0.27$.  In Fig.\\,\\ref{fig:neffvsmchi}, the Planck 68\\% and 95\\% constraints on \\neff~are shown as a function of the WIMP mass (the CMB constraints are independent of the WIMP mass).  Also shown are the curves corresponding to \\neff~as a function of \\mchi~for a Majorana fermion WIMP and for three choices of the number of equivalent neutrinos.  The behavior seen here is qualitatively similar for a Dirac or scalar WIMP (see, \\eg, \\citep{steig13} \\& \\citep{kngs}).  This figure illustrates the degeneracies between \\neff~and \\mchi.  For example, for \\Deln~= 0 the CMB can set a {\\bf lower} bound to \\mchi.  In contrast, for \\Deln~= 1 or 2, it is high values of \\mchi~that are excluded.  \\subsection{BBN Constraints}  Of the light nuclides produced during BBN, D and \\4he are the relic nuclei of choice.  To account for, or minimize, the post-BBN modifications of the primordial abundances, observations at high redshift (z) and/or low metallicity (Z) are preferred.  Deuterium (and hydrogen) is observed in high-z, low-Z, QSO absorption line systems and helium is observed in relatively low-Z, extragalactic \\hii~regions.  Even so, it may still be necessary to correct for any post-BBN nucleosynthesis that may have modified their primordial abundances.  The post-BBN evolution of D and \\4he is simple and monotonic.  As gas is cycled through stars, D is destroyed and \\4he produced.  Finally, D and \\4he provide complementary probes of the parameters (\\Deln~and \\omb) of interest.  $y_{\\rm DP} \\equiv 10^{5}{\\rm (D/H)}_{\\rm P}$ is mainly sensitive to the baryon density at BBN (\\omb) and is less sensitive to \\Deln.  In contrast, the \\4he mass fraction, \\Yp, is very insensitive to \\omb, but is quite sensitive to \\Deln.  This complementary, nearly orthogonal, dependence of D and \\Yp~on $\\eta_{10}$ and \\Deln~is illustrated in Fig.\\,\\ref{fig:YvsD}.  For the analysis here (and in \\citet{kngs}), we have adopted, $y_{\\rm DP} = 2.60 \\pm 0.12$ \\citep{pettini} and \\Yp~$= 0.254 \\pm 0.003$ \\citep{izotov}.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=0.55\\columnwidth]{hevsd13c.pdf} \\caption{BBN predicted curves of constant baryon-to-photon ratio and equivalent number of neutrinos in the \\Yp~-- $y_{\\rm DP}$ plane.  From left to right (blue), $\\eta_{10} = 7.0, 6.5, 6.0, 5.5$.  From bottom to top (red) \\Deln~= 0, 1, 2.  Also shown by the filled circle and error bars are the observationally inferred values of \\Yp~and $y_{\\rm DP}$ adopted here (see the text).} \\label{fig:YvsD} \\end{center} \\end{figure}  In contrast, the post-BBN evolution of \\3he and \\7li, the other two nuclides produced in significant abundances during BBN, is complex and model dependent.  \\3he has only been observed in the relatively metal-rich interstellar medium of the Galaxy and its BBN-predicted abundance is less sensitive to \\omb~and \\Deln~than is the BBN D abundance.  \\3he is not used in our BBN analysis, but we have confirmed that its observationally inferred primordial abundance \\citep{rood} is in good agreement with our BBN-predicted results.  \\7li suffers from some of the same issues as \\3he.  Although \\7li is observed in very metal poor stars, its post-BBN evolution is complicated and model dependent, especially the connection between the surface lithium abundances observed and those in the gas out of which the stars formed.  While in principle \\7li could be as useful as D in constraining \\omb~(and, to a lesser extent, \\Deln), there is the well known ``lithium problem\" (see, \\eg, \\citet{fields} and \\citet{spite} for recent reviews) that, as will be seen below, persists.  In the BBN analyses, with and without a light WIMP, only D and \\4he are used to constrain \\omb~and \\Deln~(or, \\neff) and these BBN constraints are compared to the independent constraints from the CMB. ",
        "conclusions": "In the absence of a light WIMP and equivalent neutrinos (no dark radiation), BBN (SBBN) depends on only one parameter, the baryon abundance.  For the adopted primordial D and \\4he abundances, SBBN predicts a best fit baryon density of $\\eta_{10} = 6.0 \\pm 0.2$ \\citep{kngs}, in excellent agreement with the corresponding value of the baryon abundance ($\\eta_{10} = 6.06 \\pm 0.07$) inferred from the Planck analysis with \\neff~fixed \\citep{planck}.  However, as shown in \\citet{kngs}, for either of these baryon abundances, BBN predicts \\Yp~$= 0.247\\,(\\pm\\, 0.0005)$, which is a poor fit to the observationally inferred primordial helium abundance of \\Yp~$= 0.254\\pm0.003$ \\citep{izotov}.  In the absence of a light WIMP, but now allowing for dark radiation (\\Deln~and \\neff~free to vary), the effective number of neutrinos and the number of equivalent neutrinos are related by \\neff~$= 3.05(1 + \\Deln/3)$.  From the Planck CMB analysis alone, \\neff~$= 3.30 \\pm 0.27$ \\citep{planck}, constraining the number of equivalent neutrinos to  \\Deln~$= 0.25 \\pm 0.27$, consistent with the absence of dark radiation (\\Deln~= 0) at $\\lsim 1\\,\\sigma$ and, inconsistent with a sterile neutrino (\\Deln~= 1) at $\\sim$\\,$2.8\\,\\sigma$.  In addition to the constraint on \\Deln, the CMB alone also provides a constraint on the universal baryon density, \\omb~$= 0.0223 \\pm 0.0003$ ($\\eta_{10} = 6.11 \\pm 0.08$).  For the Planck determined combination of the baryon abundance and the number of equivalent neutrinos, the BBN predicted D and \\4he abundances ($y_{\\rm DP} = 2.56 \\pm 0.08$ and \\Yp~$= 0.2505 \\pm 0.0005$) are in very good agreement with the observationally inferred primordial abundances, while lithium (A(Li)~$= 2.73 \\pm 0.03$) is too high.  For the same case (no light WIMP, \\Deln~free to vary), the BBN fit to the observed D and \\4he abundances is in excellent agreement with the CMB inferred parameter values (BBN: $\\eta_{10} = 6.19 \\pm 0.21$, \\Deln~$= 0.51 \\pm 0.23$, \\neff~$= 3.56 \\pm 0.23$).  A joint BBN + CMB analysis predicts $\\eta_{10} = 6.13 \\pm 0.07$ (\\omb~$= 0.0224 \\pm 0.0003$) and \\Deln~$= 0.40 \\pm 0.17$ (\\neff~$= 3.46 \\pm 0.17$).  BBN with these joint fit parameter values predicts $y_{\\rm DP} = 2.60 \\pm 0.08$ and \\Yp~$= 0.2525 \\pm 0.0005$.  The corresponding lithium abundance, A(Li)~$= 2.72 \\pm 0.03$, is a factor of $\\sim$ 3 higher than the observationally inferred, primordial value.  In the absence of a light WIMP, BBN and the CMB are in excellent agreement, but neither \\Deln~$= 0$ (SBBN) nor \\Deln~$= 1$ (a sterile neutrino) is favored by BBN or by the combined BBN + CMB fit.  In the presence of a sufficiently light WIMP (\\mchi~$\\lsim 20\\,{\\rm MeV}$) the CMB results are unchanged, although the connection between \\neff~and \\Deln~is modified depending on the WIMP mass, \\neff~$= {\\rm N}^{0}_{\\rm eff}(\\mchi)(1 + \\Deln/3)$.  Now there is a degeneracy between the CMB constraints on \\neff~and \\Deln~(and \\mchi).  As may be seen from Fig.\\,\\ref{fig:neffvsmchi}, for some choices of \\Deln, the CMB constraint on \\neff~sets a lower limit to \\mchi, while for other choices the CMB sets an upper limit to the WIMP mass.  The independent constraints from BBN help to break these degeneracies.  In the presence of a light WIMP BBN now depends on three parameters: \\Deln, \\neff, \\omb~(or, \\Deln, \\mchi, $\\eta_{10}$), but there are only two BBN constraints (from D and \\4he).  For each choice of \\mchi, corresponding to a fixed value of ${\\rm N}_{\\rm eff}^{0}$, there is always a pair of \\Deln~and $\\eta_{10}$ values for which BBN predicts -- exactly -- the observationally inferred primordial D and \\4he abundances adopted here.  However, the corresponding BBN inferred values (and ranges) of \\neff~and \\omb~need not necessarily agree with the values (and ranges) set by the CMB.  By comparing the BBN and CMB constraints, the degeneracies may be broken, leading to a lower bound, as well as a best fit value, of the WIMP mass (depending on the nature of the WIMP).  For the case of a Majorana fermion WIMP shown in the figures here, \\mchi~$\\gsim 1.7\\,{\\rm MeV}$ and the best fit is for \\mchi~$ = 7.9\\,{\\rm MeV}$.  Depending on the nature of the WIMP, the lower bound to \\mchi~ranges from $\\sim$\\,$m_{e}$ to $\\sim$\\,$10\\,m_{e}$, while the best fit WIMP masses lie in the range $\\sim$\\,$5 - 10\\,{\\rm MeV}$ (see \\citet{kngs}).  In all cases, very nearly independent of the nature of the WIMP, N$_{\\rm eff}^{0} \\approx 2.71$ and \\Deln~$\\approx 0.65$.  While the joint BBN + CMB analysis is dominated by the CMB values for \\neff~and \\omb, the presence of an additional free parameter (\\mchi) relaxes the constraints (increases the error) on \\Deln~compared to the no light WIMP case.  In the presence of a sufficiently light WIMP a sterile neutrino is now permitted at $\\lsim 68\\%$ confidence (see the right hand panel of Fig.\\,\\ref{fig:neffvsomb}).  However, the absence of dark radiation (\\Deln~= 0) remains disfavored at $\\sim 95\\%$ confidence.  For the joint BBN + CMB analysis the BBN predicted primordial lithium abundance is A(Li)~$= 2.73 \\pm 0.03$, essentially identical to that for the no light WIMP case.  The persistence of the lithium problem is largely a result of the strong coupling between the BBN predicted abundances of D and \\7li, and cannot be resolved by an extension of SBBN to include equivalent neutrinos (\\Deln~$\\neq 0$) or light WIMPs.  It should be noted that since TAUP 2013, \\citet{cooke} published new results on the primordial abundance of deuterium, $y_{\\rm DP} = 2.53 \\pm 0.04$.  Although their new central value agrees very well with the earlier, \\citet{pettini} result adopted here, the new uncertainty is smaller by a factor of three.  In the analysis described here (and, in more detail in \\citet{kngs}), this small change in the primordial deuterium abundance has the effect of increasing $\\eta_{10}$ by $\\sim$\\,$0.1$ and decreasing \\Deln~by $\\sim$\\,$0.01$.  These small changes, well within the current errors, leave the results and conclusions presented here unaffected.    \\begin{center} {\\bf Acknowledgements}\\\\ \\end{center}  We are grateful to the Ohio State University Center for Cosmology and Astro-Particle Physics for hosting K. M. N's visit during which most of the work described here was done.  K. M. N is pleased to acknowledge support from the Institute for Nuclear and Particle Physics at Ohio University.  G. S. is grateful for the hospitality provided by the Departamento de Astronomia of the Instituto Astron$\\hat{\\rm o}$mico e Geof\\' \\i sico of the Universidade de S\\~ao~Paulo, where these proceedings were written.  The research of G. S. was supported at OSU by the U.S.~DOE grant DE-FG02-91ER40690."
    },
    "1402/1402.7073_arXiv.txt": {
        "abstract": "We present the evolution of the structure of relaxed cold dark matter haloes in the cosmology from the Planck satellite. Our simulations cover 5 decades in halo mass, from dwarf galaxies to galaxy clusters.  Due to the increased matter density and power spectrum normalization the concentration mass relation in the Planck cosmology has a $\\sim 20\\%$ higher normalization at redshift $z=0$ compared to WMAP cosmology.  We confirm that CDM haloes are better described by the Einasto profile; for example, at scales near galaxy half-light radii CDM haloes have significantly steeper density profiles than implied by NFW fits.  There is a scatter of $\\sim 0.2$ dex in the Einasto shape parameter at fixed halo mass, adding further to the diversity of CDM halo profiles.  The evolution of the concentration mass relation in our simulations is not reproduced by any of the analytic models in the literature.  We thus provide a simple fitting formula that accurately describes the evolution between redshifts $z=5$ to $z=0$ for both NFW and Einasto fits. Finally, the observed concentrations and halo masses of spiral galaxies, groups and clusters of galaxies at low redshifts are in good agreement with our simulations, suggesting only mild halo response to galaxy formation on these scales. ",
        "introduction": "\\label{sec:intro}  In the standard theoretical framework for structure formation in the Universe, the mass-energy budget is dominated by a cosmological constant and cold dark matter (CDM).  In this paradigm, initially small density perturbations grow via gravitational instability, forming bound structures known as dark matter haloes.  The structure of dark matter haloes are of particular interest as they provide a non-linear scale test of the cold dark matter paradigm, cosmological parameters, and more generally for the nature of dark matter itself (e.g., Moore 1994; Flores \\& Primack 1994; de Blok \\etal 2001; Zentner \\& Bullock 2002; McGaugh 2004). They also provide the backbone for the structural properties of galaxies and galaxy scaling relations (e.g., Mo \\etal 1998; Dutton \\etal 2007, 2013).  From the theoretical side there are two hurdles that need to be overcome before an accurate prediction for the structure of CDM haloes is possible: 1) halo structure is sensitive to cosmological parameters (e.g., Macci\\`o \\etal 2008) and 2) the galaxy formation process can cause haloes to both contract or expand (e.g., Di Cintio \\etal 2014). The subject of this paper is to constrain the effects of the cosmological parameters on the ``baryon free'' dark halo structure, and to quantify the evolution of the structure of CDM haloes as a population across cosmic time.  This work continues on from our earlier studies (Macci\\`o \\etal 2007, 2008; Mu{\\~n}oz-Cuartas \\etal 2011), as well as numerous studies in the literature (e.g., Navarro, Frenk \\& White 1996, 1997; Bullock \\etal 2001; Eke \\etal 2001; Zhao \\etal 2003, 2009; Duffy \\etal 2008; Gao \\etal 2008; Klypin \\etal 2011; Prada \\etal 2012; Ludlow \\etal 2013a,b).  As shown by Macci\\`o \\etal (2008) relatively small changes in cosmological parameters have a non-negligible effect on the structure of CDM haloes. For example, the mean concentrations of CDM haloes varied by a factor of $1.5$ between the various WMAP cosmologies (Spergel \\etal 2003, 2007). The cosmology advocated by the Planck satellite (the Planck Collaboration 2013) has a significantly higher matter density, $\\Omega_{\\rm m}$, than adopted in all previous high-resolution simulations (See Fig.~\\ref{fig:omega_sigma} and Table~\\ref{tab:cosmo_param}).  Compared to the WMAP 5th year cosmology (Komatsu \\etal 2009), the Planck cosmology also has higher $\\sigma_8$. These differences are expected to result in increased dark halo concentrations (e.g., using the model of Bullock \\etal 2001).  \\begin{table*} \\centering \\caption{Cosmological Parameters. All cosmologies are flat, i.e., $\\Omega_{\\Lambda}+\\Omega_{\\rm m}=1$. } \\begin{tabular}{lcccccl} \\hline Name       &  $\\Omega_{\\rm m}$      & $h$             & $\\sigma_8$      & $n$             & $\\Omega_{\\rm b}$ & Simulation Reference\\\\ \\hline Planck     & 0.3175           & 0.671           & 0.8344          & 0.9624           & 0.0490  &  This paper \\\\ WMAP5      & 0.258\\phantom{1} & 0.72\\phantom{1} & 0.796\\phantom{1}& 0.963\\phantom{1} & 0.0438 & Macci\\`o \\etal 2008\\\\ WMAP3      & 0.238\\phantom{1} & 0.73\\phantom{1} & 0.75\\phantom{11}& 0.95\\phantom{11} & 0.042\\phantom{1} & Macci\\`o \\etal 2008\\\\ WMAP1      & 0.268\\phantom{1} & 0.71\\phantom{1} & 0.90\\phantom{11}& 1.0\\phantom{111} & 0.044\\phantom{1} & Macci\\`o \\etal 2008\\\\ Millennium & 0.25\\phantom{11} & 0.73\\phantom{1} & 0.90\\phantom{11}& 1.0\\phantom{111} & 0.045\\phantom{1} & Springel \\etal 2005\\\\ Bolshoi    & 0.27\\phantom{11} & 0.70\\phantom{1} & 0.82\\phantom{11}& 0.95\\phantom{11} & 0.0469 & Klypin \\etal 2011\\\\ \\hline \\end{tabular} \\label{tab:cosmo_param} \\end{table*}   \\begin{figure} \\psfig{figure=fig1.eps,width=0.47\\textwidth} \\caption{Constraints on the cosmic matter density ($\\Omega_{\\rm m}$) and the power spectrum normalization ($\\sigma_8$) from the Planck Collaboration (2013) (red point with 68\\% confidence intervals). For comparison, the various sets of cosmological parameters from the WMAP satellite are shown with blue hexagons (1st, 3rd and 5th year results), and the parameters used by the Millennium and Bolshoi N-body simulations are given by the black pentagon and square, respectively.} \\label{fig:omega_sigma} \\end{figure}   Traditionally the structure of CDM haloes has been described by the two parameter NFW profile (Navarro, Frenk, \\& White 1996; 1997). This has a divergent inner density profile of $\\rho(r)\\propto r^{-1}$, and an outer profile of $\\rho(r)\\propto r^{-3}$. A common parametrization uses the halo mass, $M_{200}$, and the concentration parameter, $c\\equiv r_{200}/r_{-2}$ (where $r_{200}$ is the virial radius, and $r_{-2}$ is the scale radius).  The mass and concentration are correlated, with a shallow slope ($c\\propto M_{200}^{-0.1}$), and small scatter ($\\sigma_{\\log c}\\sim 0.1$), so the structure of CDM haloes is almost scale free and described by a universal profile (NFW 97; Bullock \\etal 2001).  Recent work has shown that three parameters provide a more accurate description of spherically average CDM density profiles - especially at small radii ($\\sim 1\\%$ of the virial radius). The most common generalizations are to allow the inner logarithmic density slope, $\\gamma$, to be a free parameter (recall the NFW profile has $\\gamma=-1$), sometimes known as a generalized NFW profile (or gNFW); or the Einasto profile (Einasto 1965), which is $d\\ln\\rho / d\\ln r \\propto r^{\\alpha}$.  A number of studies have shown that the Einasto profile provides, in general, a better description of CDM haloes than the NFW or gNFW profiles (e.g., Navarro \\etal 2004, 2010; Merritt \\etal 2005, 2006; Stadel \\etal 2009; Reed \\etal 2011).  Going one step further, using stacks of CDM haloes, Gao \\etal (2008) showed that the Einasto shape parameter, $\\alpha$, depends on halo mass.  As in our earlier studies we use a large suite of cosmological N-body simulations with different box sizes to cover the entire halo mass range from $\\sim 10^{10} \\hMsun$ (haloes that host dwarf galaxies) to $\\sim 10^{15} \\hMsun $ (massive clusters).  We use these simulations to investigate the structure of CDM haloes across cosmic time. In estimating halo concentrations we consider fits to the density profiles using both the Einasto and NFW functions, as well as a non-parametric approximation utilizing $V_{\\rm max}/V_{200}$ following Klypin \\etal (2011).  This paper is organized as follows: in \\S\\ref{sec:nbody} the simulations and the determination of the halo parameters are presented.  In \\S\\ref{sec:einasto} we discuss the quality of Einasto vs NFW fits as well as different methods for measuring halo concentrations. In \\S\\ref{sec:concentration} we compare the concentration mass relation of Planck cosmology to that from WMAP, and discuss the effects of different cosmological parameters. In \\S\\ref{sec:evolution} we present the evolution of the NFW concentration mass relation, while in \\S\\ref{sec:evolution_einasto}, we present the evolution of the Einasto concentration and shape parameters. In \\S\\ref{sec:observations} we compare the concentration mass relation from our simulations with observations of spiral galaxies and clusters of galaxies.  Finally, we summarize our results in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this paper we have used a large set of cosmological N-body simulations to study the evolution of the structure of cold dark matter haloes over 90\\% of the age of the Universe. At redshift zero our simulations span five orders of magnitude in halo mass ($10^{10}-10^{15} \\hMsun$), covering haloes those that host individual dwarf galaxies to those associated with massive clusters. We adopt the cosmological parameters derived from the first data release of the Planck Satellite (the Planck Collaboration 2013). We summarize our results as follows: \\begin{itemize} \\item The concentration mass relation in the Planck cosmology has a 20\\% higher normalization (at redshift $z=0$) than in the WMAP 5th year cosmology. Despite significant differences in cosmological parameters the Planck concentration mass relation is very similar to that from the WMAP1 cosmology. By coincidence, the increased $\\Omega_{\\rm m}$ (in the Planck vs WMAP1 cosmology) is almost perfectly balanced by the decrease in $\\sigma_8, n$, and $h$. \\item Propagating the uncertainties in cosmological parameters given by the Planck Collaboration results in just a 3.5\\% uncertainty in the concentrations of Milky Way mass haloes at redshift $z=0$, which is smaller than typical systematic uncertainties in measuring halo concentrations. \\item In agreement with previous studies we find that the spherically averaged density profiles of CDM haloes are better described by the Einasto (1965) profile than the NFW (1997) profile. For example, between 2\\% and 4\\% of the virial radius, simulated haloes of mass $M_{200}\\sim10^{13}\\hMsun$ (which host massive elliptical galaxies) have average logarithmic density slopes $d\\log\\rho/d\\log r \\simeq -1.6$, compared to $\\simeq -1.4$ for NFW fits. \\item At fixed redshift the average Einasto shape parameter, $\\alpha$, increases with halo mass, from $\\alpha\\sim 0.16$ for dwarf haloes to $\\alpha\\sim 0.25$ for cluster haloes. At fixed halo mass the average, $\\alpha$, increases systematically with redshift. This evolution is well described by the relation between $\\alpha$ and dimensionless peak height $\\nu(M,z)=\\delta_{\\rm crit}(z)/\\sigma(M,z)$ proposed by Gao \\etal (2008) --- who used simulations in the WMAP1 cosmology. \\item The distribution in $\\alpha$ about the $\\alpha-\\nu$ relation is well described by a log-normal function, with a standard deviation in $\\log_{10}\\alpha$ of $\\sim 0.2$ and a slight dependence on redshift. \\item We find systematic differences of order $\\sim 10\\%$ in halo concentrations due to different fitting methods. Relative to Einasto fits to the density profile, NFW fits give concentrations that differ by up to 15\\%, while the $V_{\\rm max}/V_{200}$ method (e.g., Klypin \\etal 2011; Prada \\etal 2012) gives differences of up to 25\\%. A consequence of these systematic differences is the upturn in the concentration mass relation at high masses and redshifts (e.g., Klypin \\etal 2011; Prada \\etal 2012) is not present for our Einasto fits. \\item None of the analytic models in the literature (e.g., NFW 1997; Bullock \\etal 2001; Gao \\etal 2008; Zhao \\etal 2009; Prada \\etal 2012) accurately reproduce the evolution of the concentration mass relation. We provide simple fitting formulae for NFW (Eqs.~\\ref{eq:c200_slope}--\\ref{eq:cvir_zero}) and Einasto fits (Eqs.~\\ref{eq:cm_slope_einasto} \\& \\ref{eq:cm_zero_einasto}), which are valid between redshifts $z=5$ and $z=0$. \\item The observed concentrations and halo masses from NFW fits to data of spiral galaxies from the DiskMass project (Bershady \\etal 2010; Martinsson \\etal 2013), groups and clusters of galaxies (Ettori \\etal 2010; Auger \\etal 2013) are in good agreement with our simulations suggesting only mild halo response to galaxy formation on these scales. \\end{itemize}"
    },
    "1402/1402.2656_arXiv.txt": {
        "abstract": " ",
        "introduction": "There is evidence that long duration gamma-ray bursts (GRBs) are produced when a massive star dies at the end of its nuclear burning life and its core collapses to a neutron star or a blackhole (e.g. Galama et al. 1998, Hjorth et al. 2003, Stanek et al. 2003, Modjaz et al. 2006, Campana et al. 2006, Starling et al. 2011, Sparre et al. 2011, Melandri et al. 2012). The newly formed compact object at the center of the progenitor star produces a pair of relativistic jets that make their way out of the star along polar regions. Punching their way to the stellar surface these jets shock heat the material they encounter and push it both sideways and along the jet's direction. Therefore, the jet is surrounded by this shock heated plasma, or a hot cocoon, which provides collimation for it (see fig. 1). The total amount of thermal energy in the cocoon is equal to the work done by the jet on stellar material it encounters while inside the star and that is estimated to be of order $L_j t_*\\sim 10^{51}$erg --- where $L_j$ is the jet luminosity and $t_*$ is the travel time for the jet inside the star at sub-relativistic speed (e.g. Meszaros \\& Rees, 2001; Ramirez-Ruiz et al. 2002; Matzner, 2003). We are here making the rough assumption that the bulk of the work goes to thermal energy though a significant amount goes into the cocoon forward (in jet direction) momentum.  \\begin{figure}[ht!] \\centering \\includegraphics[scale=0.17]{cocoon-IC-sketch.eps} \\caption{Schematic sketch of a hot cocoon surrounding the jet while the jet is still inside the star is shown on the left side of this figure. The sketch to the right shows half of the system (the hemisphere with right to the center jet) at a later time when the jet and the cocoon have punched through the surface of GRB progenitor star. Thermal radiation from the cocoon is inverse-Compton scattered by the same jet, and by any relativistic jet produced by the central engine at a later time when the {\\it late jet} emerges above the cocoon surface. The IC scattered photons form a halo peaked near the edge of the jet and roughly as wide as the jet. \\label{fig1} } \\end{figure}  The jet and the cocoon emerge from the stellar surface more or less at the same time. The central engine of long-GRBs remains active for $\\sim$10s to 10$^5$s after the jet emerges from the stellar surface (e.g. Burrows et al. 2005, Chincarini et al. 2007, 2010). The isotropic equivalent of luminosity from the cocoon at its peak, shortly after it emerges from the stellar surface, is of order 10$^{49}$ erg s$^{-1}$ at a few keV (Ramirez-Ruiz et al. 2002). The cocoon luminosity decreases slowly with time for about 10$^3$s and then drops to zero rapidly when the cocoon becomes transparent to Thomson scattering of photons. Any relativistic jet that is launched after the cocoon breaks through the stellar surface is expected to encounter this intense radiation field when it reaches the cocoon photosphere, and will produce a short-lived bright pulse of inverse-Compton scattered photons which form a halo peaked near the edge of the jet and roughly as wide as the jet (fig. 1)\\footnote{The jet while inside the cocoon is shielded from the IC drag of its intense radiation by an optically thick layer of electron-positron plasma that is created at the interface of the jet and cocoon. These $e^\\pm$ pairs are produced by the collision of thermal photons from the cocoon with IC scattered photons at this interface (Ceccobello \\& Kumar, 2014).}. Furthermore, radiation from the cocoon is Thomson scattered by electrons in the wind of GRB progenitor star, and some of these wind-scattered photons collide with the relativistic jet at large angles to produce an intense pulse of high energy inverse-Compton radiation.  We provide in this work an estimate of IC flux from these different interactions of cocoon radiation with the GRB relativistic jet, and suggest how the detection of this high energy radiation (or an upper limit) can be used to constrain GRB jet and progenitor star properties.  In section 2 we describe cocoon dynamics and radiation which follows closely the work of Ramirez-Ruiz et al. (2002) and Matzner (2003), and is included here for the sake of completeness and ease to follow the rest of the paper. Section 3 describes the IC scattering of cocoon radiation by a relativistic jet directly, and via an intermediate process of scattering first by electrons in the circum-burst medium (CBM). Application to GRBs is provided in \\S4. ",
        "conclusions": "Relativistic jets of long duration GRBs push aside stellar material, and evacuate a cavity, through the progenitor star on their way out to the surface. This process creates a hot cocoon of plasma surrounding the jet with energy of order $10^{52}$ erg (isotropic equivalent). A fraction of this energy is radiated away on time scale of a few hundred seconds (in observer frame) when the cocoon punches through the stellar surface. The interaction of this cocoon radiation with the relativistic jet has been investigated in this work and shown to be useful for exploring GRB jet and progenitor star properties. The basic idea is easy to explain. The radiative luminosity of the cocoon is of order 10$^{48}$ erg/s (isotropic equivalent), and photons from the cocoon collide with the jet at an angle of order $\\Gamma_c^{-1}$ wrt jet axis; $\\Gamma_c$ is the Lorentz factor of the cocoon. The cocoon luminosity as viewed in the jet comoving frame is a factor $\\Gamma_j^2/(3\\Gamma_c^4)$ larger. The rate at which radiation is produced by the jet --- which we see as prompt $\\gamma$-ray or X-ray flare radiation --- is of order $10^{51}\\Gamma_j^{-2}$ erg s$^{-1}$ (isotropic equivalent) in its comoving frame.  Therefore, the ratio of cocoon thermal-radiation and jet radiation energy densities is $\\sim 3\\times10^{-4} (\\Gamma_j/\\Gamma_c)^4$ in the jet comoving frame. This ratio is larger than 1 for $\\Gamma_j/\\Gamma_c>10$, and in that case cocoon radiation is more important than the radiation produced within the relativistic jet for radiative cooling of electrons. Figures \\ref{ic-lum1} and \\ref{ic-lum2} show that the IC scattered cocoon radiation --- which forms a halo peaked near the edge of the jet and roughly as wide as the jet --- is of order the luminosity carried by the relativistic jet if electrons in the jet are heated to a thermal Lorentz factor larger than about 10$^2$ within a distance from the central engine of $\\sim 10^{15}$cm.  The interaction of cocoon radiation with jet and predictions for high energy emission have been investigated in detail in this work. A lack of detection of IC scattered cocoon thermal-radiation suggests either that the jet energy is not dissipated and imparted to electrons out to a radius of at least 10$^{15}$ cm --- which would rule out a certain class of models for GRB prompt emission --- or that the cocoon moves outward with a high Lorentz factor such that $\\Gamma_j/\\Gamma_c \\lae 5$ (this possibility can be constrained by afterglow observations).  Photons from the cocoon scattered by electrons in the circum-burst medium (CBM) can collide with the jet at a larger angle than photons traveling from the cocoon to the jet directly. These collisions result is very high energy photons ($\\sim$GeV) of considerable luminosity even for a modest thermal Lorentz factor for electrons (fig. \\ref{ic-lum2}). There is considerable uncertainty, however, in this estimate because the CBM could have been partially evacuated by an earlier passage of a relativistic jet through this region. Detection of this signal, or an upper limit, would provide a handle on the stellar mass loss rate during the last few years of the life of the GRB progenitor star.  \\medskip \\noindent{\\bf Acknowledgment:} The work of GS is funded in part by PCCP. \"Paris Center for Cosmological Physics acknowledges the financial support of the UnivEarthS Labex program at Sorbonne Paris Cit\u00e9 (ANR-10-LABX-0023 and ANR-11-IDEX-0005-02)\"."
    },
    "1402/1402.5240_arXiv.txt": {
        "abstract": "Recent observational and theoretical studies of classical Be stars have established the utility of polarization color diagrams (PCD) in helping to constrain the time-dependent mass decretion rates of these systems.  We expand on our pilot observational study of this phenomenon, and report the detailed analysis of a long-term (1989-2004) spectropolarimetric survey of 9 additional classical Be stars, including systems exhibiting evidence of partial disk-loss/disk-growth episodes as well as systems exhibiting long-term stable disks.   After carefully characterizing and removing the interstellar polarization along the line of sight to each of these targets, we analyze their intrinsic polarization behavior.  We find that many steady-state Be disks pause at the top of the PCD, as predicted by theory.  We also observe sharp declines in the Balmer jump polarization for later spectral type, near edge-on steady-state disks, again as recently predicted by theory, likely caused when the base density of the disk is very high, and the outer region of the edge-on disk starts to self absorb a significant number of Balmer jump photons.  The intrinsic $V$-band polarization and polarization position angle of $\\gamma$ Cas exhibits variations that seem to phase with the orbital period of a known one-armed density structure in this disk, similar to the theoretical predictions of Halonen \\& Jones.  We also observe stochastic jumps in the intrinsic polarization across the Balmer jump of several known Be+sdO systems, and speculate that the thermal inflation of part of the outer region of these disks could be responsible for producing this observational phenomenon.  Finally, we estimate the base densities of this sample of stars to be between $\\approx 8\\times 10^{-11}$ to $\\approx 4 \\times 10^{-12}\\,\\rm g cm^{-3}$ during quasi steady state periods given there maximum observed polarization. ",
        "introduction": "\\label{intro}  Classical Be stars are a subset of B-type main sequence stars which are characterized by their rapid rotational velocities ranging from 60\\% to 100\\% of their critical rate \\citep{riv13}.  They have a geometrically flattened decretion disk that is fed from material from the stellar photosphere as diagnosed from studies of their optical/IR emission lines, polarization, and interferometric signatures (see e.g. \\citealt{por03,stee11}).  A large volume of observations suggest the kinematic properties of these gas disks is best represented by near Keplerian rotation \\citep{hum00, me07a, pot10, whe12, kra12}.  For the most up to date review of Classical Be stars see \\cite{riv13}.  As summarized in \\citet{car11}, the viscous decretion disk model developed by \\citet{lee91} can explain many of the observational signatures of Be disks, although other models have been explored to explain the structure of these disks \\citep{bjo93,cas02,bro08}.  One key unanswered question in the study of Be disks is what mechanism(s) are responsible for injecting material into these disks.  Non-radial pulsations have been suggested to be one contributing factor to supplying material to some of these disks \\citep{cra09,riv98,nei02}, while periastron passage of binary companions may contribute in other systems such as $\\delta$ Scorpii \\citep{mir01,mir03}.  This scenario for $\\delta$ Scorpii is now questionable given the disk's growth prior to the periastron passage of 2011 \\citep{miro13}.  Nonetheless, binarity may play a role in the phenomenon and disk variability.  For example, the source of material and angular momentum in non-classical Be stars can be the result of a red giant phase binary transferring material to create a Be+sdO system \\cite{gies98}. Characterizing the evolution of Be stars' mass-loss rates is another promising approach to constrain the disk-feeding mechanism.  On short time-scales, \\cite{car07} noted polarimetric variability in Archernar likely arising from injections of discrete blobs of mass into the inner disk, that subsequently circularize into rings.  Studying longer duration disk-loss and disk-regeneration events \\citep{und82,doa83,cla03,vin06,hau12}, including the time-scales \\citep{wis10} and statistical frequency \\citep{mcs08,mcs09} of these episodes, is another way to diagnose the mechanism feeding Be disks.  Polarimetry has been used to study the Be phenomenon for both individual Be stars \\citep{qui97,woo97,cla98} and larger statistical surveys \\citep{coy69,mcl78,poe79,wi07b}.  It is widely believed that Thompson scattering (free-electron scattering) is the source of polarization in Be stars \\citep{wo96a,wo96b,hal13}.  Pre- or post-scattering absorption of photons within the disk can imprint wavelength dependent signature on top of the wavelength independent Thompson scattering \\citep{woo95}.  Because the polarization across the Balmer jump traces material in the innermost regions of disks ($\\sim$6 R$_{\\star}$; \\citealt{car11}, \\citealt{ha13a}) and the $V$-band polarization is a tracer of the total scattering mass of the disk, studying the time evolution of the wavelength-dependence of polarization in Be stars can be used to constrain the time dependence of the mass decretion rate and the $\\alpha$ parameter in these systems.  Specifically, the slope, shape, and temporal evolution of polarization color diagrams (PCD) observed in Be systems \\citep{dra11} have been theoretically reproduced, for the first time, by time-dependent radiative transfer models in which the mass decretion rate, $\\alpha$ parameter, and inclination angle are the primary variables \\citep{hau13}.  In addition to density changes caused by changes in the mass decretion rate, it has been suggested that PCD diagram loops can also be produced by one-armed density perturbations \\citep{ha13b}.  In this paper, we implement the PCD diagram diagnostic developed in \\citet{dra11} on a broader sample of Be stars, including systems showing evidence of experiencing disk-loss events and those whose disks appear roughly stable over time.  We describe the data sample in Section 2, and discuss the techniques we used to remove the interstellar polarization component from each dataset in Section 3.  We discuss the temporal evolution of each Be disk in PCD diagram parameter space in Section 4, and also detail evidence of variability in the disk position angle in select systems.   Finally, we summary the major results of this manuscript in Section 5. ",
        "conclusions": "\\label{discussion}  \\subsection{Additional Disk-Growth and Loss Phases}  In paper I of this series, we analyzed well sampled, long-term spectropolarimetric data for two classical Be stars that clearly underwent a complete disk-loss phase.  In this paper, we explore the behavior of a larger sample of classical Be stars, including systems that exhibit disks that are stable over time periods of more than a decade, as well as systems that exhibit at least partial disk-loss/disk-renewal phases.  While detailed modeling of each system that exhibits a non-stable disk is clearly necessary and warranted, albeit outside the scope of this paper, we discuss some of the basic observational properties of these systems below.  \\subsubsection{$\\psi$ Per}  The intrinsic polarization and H$\\alpha$ EW of $\\psi$ Per both exhibit clear evidence of significant variability over the $\\sim$13 year time coverage of our data (Fig. \\ref{fig:ip_all_b}).  Throughout the early 1990s, the steady increase in the system's intrinsic polarization, lasting for $\\sim$1400 days, accompanied by a small strengthening in its H$\\alpha$ EW is indicative of a growth in density and size of the disk.  Although the data sampling is sparse, there is evidence that the disk generally stabilized in strength between 1995-2001.  Figure \\ref{fig:ip_all_b} also exhibits a dramatic, monotonic drop in intrinsic polarization over a period of $\\sim$190 days starting in August 2002, along with a more gradual, time-delayed, significant decrease in H$\\alpha$ EW.  These trends are consistent with a major inside-out clearing of significant mass from the disk \\citep{wis10}.  Since a small level of both intrinsic polarization and H$\\alpha$ emission remained after the event, we characterize this as a incomplete disk-loss event.  Characterizing the time-scale of disk-loss events is important in order to constrain the viscosity parameter, $\\alpha$, assuming the disk dissipates on a viscous time-scale.  Equation (19) of \\citet{bjo05} describes such an assumption and can be rewritten to give $\\alpha$ as function of the diffusion timescale, $t_{\\rm diff}$.  \\begin{equation} \\alpha = (0.2 \\rm yr / t_{\\rm diff}) * (r/R_{\\star})^{0.5} \\end{equation}  Assuming the bulk of the scattering events producing the observed polarization occurs at a radial distance of $\\sim$5 R$_{\\star}$ and the (partial) disk-loss time-scale of $\\sim$ 190 days yields $\\alpha$ $\\sim$0.86 for $\\psi$ Per.  This value is intermediate compared to estimates for 60 Cyg ($\\sim$0.1; \\citealt{wis10}) and robustly measured for 28 CMa (1 $\\pm$ 0.2; \\citealt{car12}).  Because the $\\psi$ Per disk did not completely clear out, our quoted disk-loss time-scale is a lower limit and correspondingly the quoted $\\alpha$ parameter is an upper limit.  Clearly, characterizing the viscosity parameter for a larger number of systems is required to assess whether there is a preferred parameter for most Be disks.  \\subsubsection{$\\omega$ Ori}  The overall steady decline in H$\\alpha$ EW strength throughout most of the 13 years of coverage in our dataset (Fig. \\ref{fig:ip_all_b}) is indicative of a gradual, albeit incomplete, loss of the system's disk.  The $\\sim$1500 day duration decline in intrinsic $V$-band polarization starting at the beginning of our dataset is consistent with a gradual loss of the system's disk.  Mirroring the scenario observed in $\\pi$ Aqr reported by \\citet{wis10}, the subsequent $\\sim$1100 day long increase in intrinsic polarization is likely responsible for the temporary halt in the decline of H$\\alpha$ EW strength between 1995-1998 and indicative of temporary replenishment of material to the inner regions of the disk.  The cessation of this mass injection to the inner disk is marked by a fast decrease in the intrinsic polarization level, as well as a resumption of the gradual decline in H$\\alpha$ EW.  Near the end of our dataset, the sharp rise in both intrinsic polarization and H$\\alpha$ EW indicate significant mass injection into the disk.  Overall, these disk-loss and growth epochs occurred in the span of 2-4 years.  \\subsubsection{66 Oph} The H$\\alpha$ EW of 66 Oph also exhibited a steady decline throughout the entire 13 year baseline of our dataset (Fig. \\ref{fig:ip_all_b}).  Supplementary H$\\alpha$ EWs from spectra obtained at Ritter Observatory improve the time sampling during the end of our spectropolarimetric data, and confirm that the emission strength decline continues during these epochs.  The Ritter EWs end their decline after the last epoch of our spectropolarimetric data (see e.g. Table \\ref{66Oph_rit_data}), suggesting the system's mass-loss rate changed and a complete disk-loss phase was avoided.  The intrinsic polarization data (Fig \\ref{fig:ip_all_b}) also exhibit a slow decline throughout the 13 year baseline of the HPOL data, supporting the interpretation of a gradual decline in the disk.  \\subsection{Intrinsic Polarization Color Diagrams} \\label{intpol}  \\citet{dra11} noted the evolution of the ratio of the polarization across the Balmer jump versus the $V$-band polarization occasionally traced out distinctive loops in PCD diagrams. \\citet{hau13} explored these diagrams in detail with theoretical models.  Since the polarization change across the Balmer jump is roughly proportional to the density squared whereas the $V$-band polarization is proportional to the density, PCD diagrams offer a useful method to investigate the evolution of Be disks when spectral type and inclination can be constrained \\citep{hau13}.  PCD loops observed in 60 Cyg and $\\pi$ Aqr were qualitatively modeled using the radiative transfer code HDUST and 1D hydrodynamical code SINGLEBE by turning the mass decretion rate on and off.  Subsequent detailed modeling of PCDs demonstrated that the shape, slope, and time-scale of loops depends on spectral type, inclination angle, the base density of the disk, and the temporal behavior of the mass decretion rate \\citep{hau13}.   One important conclusion from these modeling efforts is that steady state disks pause at the top of a PCD loop, while the disk is fed by either constant mass-loss or semi-regular mass injections from the star \\citep{hau13}.  Our steady disk sample broadly reproduces this behavior.  For example, both $\\phi$ Per and $\\gamma$ Cas (Fig \\ref{fig:ip_all}) cluster at high $V$-band polarization and high Balmer jump ratios in their PCDs during their steady-state disk stage.\\\\  However, the following systems have deviations from this steady state disk behavior:  \\\\  \\subsubsection{48 Lib}  48 Lib exhibits clear evidence of a partial PCD loop (Fig \\ref{fig:ip_all_a}) in our dataset.   The rise in 48 Lib's intrinsic $V$-band polarization (blue points; Fig \\ref{fig:ip_all_a}) indicates a disk growth phase, which eventually plateaus into a steady state disk phase. Unlike other steady state disks in our sample that sit at the top of the PCD loop during the steady state, 48 Lib's Balmer jump ratio rapidly drops during this phase. Since 48 Lib is a near edge-on inclination \\citep{riv06,ste12}, its PCD behavior seems consistent with the modeling scenario outlined in Fig. 9 of \\citet{hau13}.  Specifically, these results suggest that 48 Lib was characterized by a high base density ($\\approx$ 10$^{-11}$ g cm$^{-3}$) disk and disk orientation which gives a high opacity.   This manifests during the onset of the disk-growth event with a large Balmer jump ratio which later decreases at a nearly constant $V$-band polarization as it approaches a steady state phase. Furthermore, it has a steeper decline which is consistent with a later spectral type of B2 to B5 given 48 Lib is B3. \\\\  \\subsubsection{$\\psi$ Per} $\\psi$ Per experienced a notable rise in intrinsic $V$-band polarization during the early 1990's (blue points; Fig \\ref{fig:ip_all_b}) indicating a disk growth phase, that seemingly plateaued during the later part of the decade into a stable phase.  $\\psi$ Per exhibited a dramatic drop in its Balmer jump polarization as it transitioned to a steady state system.  Given the mid spectral type (B5) and near edge-on inclination ($\\sim$75$^{\\circ}$; Table 1) of the system, like 48 Lib, we suggest that the $\\psi$ Per PCD behavior could be caused by the outer regions of the high inclination, high density disk absorbing significant Balmer jump photons at the onset of the disk growth event \\citep{hau13}.  This is further supported by the fact the polarization blueward of the Balmer jump was nearly zero, or essentially optically thick, to polarization of the disk. \\\\  \\subsubsection{28 Cyg} During 28 Cyg's overall long-term decline in disk strength, as measured by its declining intrinsic $V$-band polarization and H$\\alpha$ EW (Fig. \\ref{fig:ip_all_c}), the system experienced three short, large jumps in Balmer jump polarization.  During a 90 day period when the star was monitored regularly, in some cases nightly, the PCD diagram illustrates 3 instances where the Balmer jump ratio spiked to $>$2.5 (compared to the median value of 1.4). The separation between these three peaks was 24 and 18 days respectively.  Given its spectral type (B2.5), the Balmer jump polarization is most sensitive to the inner regions of the disk when compared to the $V$-band polarization \\citep{hau13}. These Balmer jump ratio spikes (that occur at a constant $V$-band polarization) are likely caused by discrete events in the inner disk.  Specifically, we speculate that the spikes stem from stochastic mass injections into the disk. \\\\  Although outside the scope of this paper, it is clear that detailed modeling of 48 Lib, $\\psi$ Per, and 28 Cyg using codes such as \\citet{ha13a} and \\citet{hau13} should be improved and pursued to reproduce the PCD variability observed in these systems. Moreover, given the diagnostic potential for systems experiencing full disk growth/disk-loss episodes ($\\pi$ Aqr, 60 Cyg; \\citealt{dra11}) and episodic disk growth (48 Lib, $\\psi$ Per, 28 Cyg), it is certainly clear that enhanced observational monitoring of these types of systems should be aggressively pursued.  Without consistent temporal monitoring, it can become challenging to constrain the exact time-scale for disk growth and loss events.  Our results for 28 Cyg provide quantitative evidence that even high cadence (i.e. nightly) observations yield interesting PCD phenomenon that could be better exploited to diagnose episodic mass injection events. \\\\  \\subsection{One-armed Density Waves} \\label{PAvar}  \\citet{ha13b} presented ad hoc model predictions of the effects of global one-armed oscillations \\citep{oka91,oka97} on the time-dependent linear polarization of Be disks, including models that used a perturbation pattern characterized by a pattern of opposite overdense and underdense regions \\citep{oka97} and models that used a spiral shaped perturbation pattern as adopted by \\citet{car09}.  These models (and those of \\citealt{car09}) predict that the $V$-band polarization and polarization across the Balmer jump should exhibit a clear inclination-dependent modulation with phase (Figs 6-8 of \\citealt{ha13b}), although they will be out of phase with one another due to the different radial locations in the disks over which the scattering events occur in each bandpass.  The $V$-band polarization position angle is also predicted to exhibit complex changes with phase (Fig 10; \\citealt{ha13b}).  In-spite of these predictions, no conclusive observational evidence of this phenomenon has been reported \\citep{car09}.  $\\gamma$ Cas is known to exhibit V/R variations indicating the presence of a one-armed density feature in its disk, as seen in its phase-folded He I V/R ratios (Fig \\ref{wacky}) compiled from data presented in \\citet{mir02}.  Observations were made at the Ritter observatory from 1993 to 2002 of the He I at 5876 \\AA, amongst other lines, with a resolving power of $\\sim$26000.  The line profile was consistently double peaked and had clear V/R variations  $\\gamma$ Cas exhibits a generally stable disk, as diagnosed by its intrinsic polarization (Fig. \\ref{fig:ip_all_a}), which indicates its inner disk is being supplied by a generally stable decretion rate. It therefore makes an ideal system to search for long-term polarimetric effects related to its one-armed density feature.  The intrinsic $V$-band polarization, polarization across the Balmer jump, and $V$-band polarization position angle phased to the V/R period of the He I data is shown in Fig \\ref{wacky}.  The intrinsic $V$-band polarization generally exhibits a double-oscillation pattern over one period that is predicted by \\citet{ha13b}. The intrinsic polarization across the Balmer jump exhibits variability, albeit no clear indication of the phase-lagged behavior predicted by \\citet{ha13b}.  The intrinsic $V$-band polarization position angle suggestively exhibits evidence of cohesive variations as a function of phase, but such variations are an order of magnitude greater than that predicted in the ad hoc models of \\citet{ha13b}.  We speculate that one reason these data exhibit a stronger indication of the predicted behavior in $V$-band, rather in the Balmer jump polarization, is that the former is less sensitive to small changes in the inner disk caused by changes in the mass decretion rate.  We encourage future modeling efforts of this phenomenon to explore the ramifications of abrupt and gradual changes in the mass decretion rate have on the polarization in disks having one-armed density waves.  \\subsection{PA Deviations}  \\citet{wis10} reported numerous instances of the intrinsic polarization of 60 Cyg and $\\pi$ Aqr deviating from their linear trends on a Stokes QU diagram, and found that these deviations in intrinsic polarization position angle were more prominent during large outburst events.  These authors interpreted this behavior as either evidence of the injection and subsequent circularization of new blobs of mass into the inner disk region, similar to that noted in \\citet{car07}, or as evidence of the injection and subsequent circularization of new blobs at an inclined orbit to the plane of pre-existing disk material. Many of the Be systems explored in our current study exhibit generally stable, strong disks over most of the duration of our dataset, which suggests we should see analogous evidence of PA deviations in our dataset.     As seen in Fig. \\ref{fig:gamCasdev}, $\\gamma$ Cas exhibits PA deviance as a function of the intrinsic $V$-band polarization in the system. Analogous figures for our other targets are available in the online-version of this paper. Overall, we do see clear evidence of strong PA deviations in most of our Be stars even though they cannot be correlated to specific outburst events like 60 Cyg and $\\pi$ Aqr were. The system that exhibits the smallest level of PA deviations, $\\omega$ Ori, is noteworthy as it exhibited evidence of a gradual, albeit incomplete, loss of its disk throughout the time-frame covered by our data.  While errors in the ISP estimate could be inducing errors in the intrinsic PA, we note that the deviations are symmetric about the mean.  If there were errors in our PA determinations, then one would expect the deviations to be systematically offset to one side of the mean.  If there were a PA estimate error then one would expect the deviations systematically offset to one side of the mean.  In general, our results are consistent with systems which have had more precise ISP determinations and time resolved outbursts which suggest the PA deviations are also ``clumpy'' injection events (e.g. 60 Cyg and $\\pi$ Aqr).   \\subsection{Be+sdO systems}\\label{sdO}  Three of our targets are known to have a sub dwarf companion, $\\phi$ Per \\citep{gies98}, FY CMa \\citep{peters08}, and 59 Cyg \\citep{peters13}, and all three systems exhibit steady state disks in our data.  This apparent stability could be influenced by the sdO truncating these disks to the maximum allowed radius.  Interestingly, we do observe significant variability in the polarization across the Balmer jump in both FY CMa and 59 Cyg (Fig \\ref{fig:ip_all_c}).  We remind the reader that the ISP for 59 Cyg was poorly constrained; however, we explored different assumed total ISP PA values (about PA = 0) and still observed the jumps in the polarization across the Balmer jump, suggesting that they are likely real.  As the sdO is likely heating the outer region of the disk nearest its orbital position, we speculate that this outside heating source may inflate the scale height of this region of the disk.  Given the short periods of order 29-37 days \\citep{peters13}, we speculate that when the inflated scale heights of these moderately high inclination disks (Table 1) pass in front of our line of sight with the star, they cause the observed changes in polarization across the Balmer jump.   We further speculate that the amount of disk material heated and inflated by $\\phi$ Per's more distant (127 day period) sdO companion may contribute to why this system exhibits no analogous Balmer jump polarization jumps.  Higher cadence observations that sufficiently sample the suggested puffed-up outer disk over the orbital period of these systems' sdO companion would help to further establish our interpretation of this observational phenomenon.  \\subsection{Maximum Polarization}  The maximum polarization for a viscous disk in equilibrium is found to be correlated with inclination, spectral type, and base density \\citep{hau13,ha13b}.  In most cases the disks in this sample reach some form of steady state.  We then assume that the maximum observed polarization at some point in the 15 year survey is the maximum polarization of a disk around these stars.  Since the effective temperature and inclination can be determined by other means (e.g. spectroscopy and interferometry), we use literature values to then derive the range of base densities for this sample of Be stars (See Table \\ref{star_sum}).  Given the model tracks of \\cite{hau13}, we find that the base density for these stars lie within $ 8\\times 10^{-11}$ to $ 4 \\times 10^{-12}\\,\\rm g cm^{-3}$ (See Fig.\\ref{pmax}).  Several items are key to interpreting the base density estimates extracted from Figure 5.  $\\pi$ Aqr, for example, exhibits a lower maximum V-band polarization than the 4.2e-11 g cm${-3}$ base density track compared to other stars having a similar stellar effective temperature (Figure 5; right panel) due in part because the system's inclination angle (33.6$^{\\circ}$ is lower than that used for the model track (70$^{\\circ}$.  Similarly, $\\pi$ Aqr  exhibits a higher maximum than expected for the 4.2e-11 g cm${-3}$ base density track (Figure 5; left panel) compared to other systems having a similar inclination due in part to $\\pi$ Aqr having a warmer stellar effective temperature than adopted for these models.  Due to this degeneracy, it likely has a similarly large base density as $\\phi$ Per around $4.2 \\times 10^{-11} \\rm g cm^{-3}$.  $\\psi$ Per and 59 Cyg have similar maximum polarization and inclination yet have very different effective temperatures.  This then requires a low base density of $8.4 \\times 10^{-12} \\rm g cm^{-3}$ which exhibits maximum polarization independent of effective temperature.  An object like $\\gamma$ Cas, shows behavior similar to that of $\\pi$ Aqr in relation to the model tracks. It is likely limited in maximum polarization due to its low inclination rather then its spectral type, so it is consistent with a low base density of $4.2 \\times 10^{-11} \\rm g cm^{-3}$.  Due to the degeneracy and variable nature of some of the stars observed by HPOL, these limits are not meant to be absolute but rather a first look into the disk properties for the class as a whole.  Each star will require more detailed modeling, but potentially broader statistics could be applied to a wider sample if polarimetry, spectroscopy, and interferometry can be obtained simultaneously to derive the applicable parameters."
    },
    "1402/1402.0379_arXiv.txt": {
        "abstract": "{Low-mass X-ray binaries (LMXBs) are a natural workbench to study accretion disk phenomena and optimal background sources to measure elemental abundances in the Interstellar medium (ISM). In high-resolution XMM-\\textit{Newton} spectra, the LMXB SAX~J1808.4$-$3658 showed in the past a neon column density significantly higher than expected given its small distance, presumably due to additional absorption from a neon-rich circumstellar medium (CSM).} {It is possible to detect intrinsic absorption from the CSM by evidence of Keplerian motions or outflows. For this purpose, we use a recent, deep (100\\,ks long), high-resolution \\textit{Chandra}/LETGS spectrum of SAX~J1808.4$-$3658 in combination with archival data.} {We estimated the column densities of the different absorbers through the study of their absorption lines. We used both empirical and physical models involving photo- and collisional-ionization in order to determine the nature of the absorbers.} {The abundances of the cold interstellar gas match the solar values as expected given the proximity of the X-ray source. For the first time in this source, we detected neon and oxygen blueshifted absorption lines that can be well modeled with outflowing photoionized gas. The wind is neon rich (Ne/O\\,$\\gtrsim$\\,3) and may originate from processed, ionized gas near the accretion disk or its corona. The kinematics ($v=500-1000$\\,km\\,s$^{-1}$) are indeed similar to those seen in other accretion disks. We also discovered a system of emission lines with very high Doppler velocities ($v\\sim24\\,000$\\,km\\,s$^{-1}$) originating presumably closer to the compact object. Additional observations and UV coverage are needed to accurately determine the wind abundances and its ionization structure.} {} ",
        "introduction": "\\label{sec:introduction}  Accretion disks provide an important workbench to probe magnetized plasma dynamics, photoionization, thermal and ionization equilibria. The response of the ionized plasma to changes in the continuum and outflow kinematics may constrain the source that drives winds, the connection between jets, disk, and corona, and possibly the disk geometry. Low-mass X-ray binaries (LMXBs) consist of a neutron star (NS) or black hole (BH) in orbit with a $<$1 M$_{\\odot}$ companion. These sources are characterized by X-ray emission due to hot gas accreted by the compact objects in a form of disks. Some of the NS binaries exhibit X-ray bursts when enough of this material accumulates on the surface of the NS, eventually giving rise to a runaway thermonuclear explosion. Bursts can last up to an hour, but they provide only a few\\,\\% of the total fluence (see e.g. \\citealt{Lewin1993}).  LMXBs display high flux and simple spectra, whose continuum is typically well described with blackbody (BB) and powerlaw (PL) emission. The BB component arises from the thermal emission of the accretion disk around the neutron star. The powerlaw describes the energy gain of the disk soft photons by scattering in the accretion disk corona, a process known as Comptonization. A few sharp features are seen in the high-energy part of the spectra, e.g. above 1\\,keV. Prominent Fe\\,K emission lines are commonly seen in LMXB high-quality spectra and a clear asymmetry is seen in the line profile, as would be expected if the lines originate from the innermost region of the accretion disk and therefore are subject to strong relativistic effects (see e.g. \\citealt{Cackett2012} and references therein).  Most LMXB spectra do not exhibit intrinsic sharp absorption features in the soft X-ray energy band, but they are rich with absorption lines originating in the interstellar medium (ISM) that lies along their line-of-sight (LOS). The ISM influences the Galactic evolution through the exchange of matter with the stars and shows a complex structure consisting of phases at different equilibrium temperatures \\citep[for a review, see][]{Draine2011}. The K-shell transitions of carbon, nitrogen, oxygen, neon, and magnesium, and the L-shell transitions of iron fall in the soft X-ray energy band. The launch of the {XMM-\\textit{Newton}} and \\textit{Chandra} satellites, provided with high spectral resolution gratings, permitted to determine interstellar ionization states and amounts of dust in the LOS towards several X-ray sources (see e.g. \\citealt{paerels, devries, JuettI, costantini2005, Costantini2012, costantini, kaastra09, Lee2009, Pinto2010}).  Recently, \\citet{Pinto2013}, hereinafter P13, have performed an extended analysis of the ISM towards nine LMXBs with high-quality {XMM-\\textit{Newton}} spectra, extracted total abundances, dust depletion factors, and attempted to constrain dust chemistry. They also measured the abundance trends in the Galactic disk and confirmed the well known metallicity gradient (see e.g. \\citealt{GradPedicelli}). However, their interstellar abundances significantly deviate from a monotonic distribution. The extreme case is given by the LMXB \\object{SAX~J1808.4$-$3658} (hereinafter SAX~J1808) where the neon abundance exceeds by 5$\\sigma$ the value predicted by the source location. P13 argue that this large scatter may be due to absorption by metal-rich material surrounding the X-ray sources.  Absorption occurring near the source may be distinguished from Galactic absorption when velocity shifts or Doppler broadening are observed. However, the XMM-\\textit{Newton} spectral resolution may not be high enough to resolve them. Recent work by \\citet{Schulz2010} on the high-resolution \\textit{Chandra} spectra of LMXB 4U\\,0614$+$091 showed a variable Ne\\,K edge with an average velocity smear of $\\sim3500$\\,km\\,s$^{-1}$ implying a characteristic radius $<10^7$\\,m, consistent with an ultra-compact binary (UCB) nature. The variability proves that the excess is intrinsic to the source, but it is not yet clear whether this is due to either neon~/~oxygen absorption or some other process. \\citet{Costantini2012} found an absorption-line system in the RGS spectrum of the LMXB 4U1820--30, which may also be described as a $\\sim1200$\\,km\\,s$^{-1}$ outflow. \\citet{Ioannou2003} found evidence of a $\\sim1500$\\,km\\,s$^{-1}$ \\ion{C}{iv} outflow in the \\textit{Hubble Space Telescope} (HST) UV spectrum of the LMXB X2127+119.  The X-ray transient SAX~J1808 is the first discovered accretion-powered millisecond X-ray pulsar in a LMXB (\\citealt{zand1998, Wijnands1998}). It is relatively nearby ($d=3.5\\pm0.1$\\,kpc, see \\citealt{Galloway2006}) and has a column density of interstellar medium $N_{\\rm H} = (1.40\\pm0.03)\\,\\times 10^{25}$\\,m$^{-2}$ (see P13). This $N_{\\rm H}$ value combined with the high source flux produces several strong interstellar absorption features in soft X-ray spectra. The source exhibits X-ray bursts showing photospheric expansion (see e.g. \\citealt{Galloway2008}). These most likely originate in a flash of a pure helium layer that is produced by stable hydrogen burning. This source is thus particularly suitable to study X-ray bursts in great detail and in a large bandpass.  Therefore, \\citet{zand2013}, hereinafter ZA13, observed SAX~J1808 with Chandra and RXTE during the November 2011 outburst. They detected a single thermonuclear (type-I) burst, the brightest yet observed by Chandra from any source, and the second-brightest observed by RXTE. However, they found no evidence for discrete spectral features during the burst $-$ though absorption edges have been predicted to be present in such bursts $-$ and argued that a greater degree of photospheric expansion may be required. The persistent spectrum is instead optimal to study the absorption edges produced by diffuse (ISM) and circumstellar (CSM) media.  In this work we used the deep observation of SAX~J1808 presented in ZA13, which was taken with the Advanced CCD Imaging Spectrometer (ACIS) in combination with the Low-Energy Transmission Grating (LETG) on board of {\\textit{Chandra}} during a high-flux source state as well as the archival observations taken with the {\\textit{Chandra}} Medium-Energy Grating (MEG) and the {XMM-\\textit{Newton}} Reflection Grating Spectrometer (RGS).  The paper is organized as follows. In Sect.~\\ref{sec:data} we present the data reduction. In Sect.~\\ref{sec:spectra} we describe the spectral features and the fitting procedure. Alternative methods of analysis are treated in detail in Sect.~\\ref{sec:spectral_models1} and \\ref{sec:spectral_models2}. In Sec.~\\ref{sec:variability} we analyze the spectral variability. All results and their comparison with previous work are discussed in Sect.~\\ref{sec:discussion}. Conclusions are given in Sect.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have presented a detailed analysis of the soft X-ray spectra at different epochs of the LMXB SAX\\,J1808.4-3658. We have carefully measured the abundances of the ISM which lies in the LOS towards this source. They all agree with the abundances of the Solar System as expected due to the nearness of the source.  In this paper we took into account not only the cold ISM phase, but we also looked for absorption by gas in the source system. A better modeling of the absorption edges region allowed us a precise estimates of elemental abundances. We found an ISM Ne abundance consistent with the Solar value, which is less than previously estimated for this source.  We detected gas outflowing at $500-1000$\\,km\\,s$^{-1}$ as shown by strong, blueshifted absorption lines at different ionization states that are well described with photoionized absorbers. We ascribed the outflow to the photoionized winds that have been observed in several sources (see e.g. \\citealt{DiazTrigo2012}). These winds are thought to originate from thermal and radiation pressure near the accretion disk. Our velocity estimates are in agreement with the dynamics of the accretion disks, but our ionization parameters are slightly lower than usual. However, we notice that outflows of intermediate-ionization gas have also been observed.  The Ne/O abundance ratio in the wind is supersolar which would suggest high Ne/O abundances for the donor star if the wind originates from the accretion disk. This would disagree with the hypothesis of the companion to be a brown dwarf unless the matter is expelled from the stellar interior.  {We also detected a system of emission lines in the spectrum of the source after a very powerful burst was exhibited. These lines can be well modeled with collisionally-ionized gas,  originating presumably closer to the compact object, but they require very high Doppler velocities ($v\\sim0.1c$). Additional, higher-resolution X-ray (MEG) and far-UV (HST) observations covering the $1-20$\\,{\\AA} and $1000-2000$\\,{\\AA} spectral ranges can improve the abundance estimates of the winds and the interpretation of the emitting plasma.}  Our method can be easily tested in several X-ray binaries in order to search for wind signatures and to probe the local circumstellar and the diffuse interstellar media."
    },
    "1402/1402.0523_arXiv.txt": {
        "abstract": "Each of the two pulsars in the double pulsar PSR J0737-3039A/B system exhibits not only the pulses emanating from itself, but  also displays modulations near the pulse period of the other. Freire et al. (2009, MNRAS, 396, 1764) have put forward a technique using the modulation of B by A to determine the sense of rotation of pulsar A relative to its orbital motion, among other quantities.  In this paper, we present another technique with the same purpose. While the Freire et al. approach analyzes pulse arrival times, ours instead uses periods or frequencies (their inverses), which can be experimentally determined via power spectral analysis similar to that used in pulsar searches.  Our technique is based on the apparent change in spin period of a body when it is measured from an orbiting platform (the other pulsar), and is shown to be entirely analogous to the difference between the sidereal and solar spin period of the Earth (i.e., the  sidereal and solar day).  Two benefits of this approach are its conceptual and computational simplicity.  The direct detection of spin with this technique will observationally validate the rotating lighthouse model of pulsar emission, while the detection of the relative directions of spin and orbital angular momenta has important evolutionary implications. Our technique can be used on other binary systems exhibiting mutually induced phenomena. ",
        "introduction": "The double pulsar system PSR J0737-3039A/B was discovered by \\citet{Burgay03} and \\citet{Lyne04}. This system consists of a 22-ms pulsar (hereafter A) and a 2.8-s pulsar (hereafter B) with an orbital period of 2.4 hours. This discovery has provided a laboratory for the study of relativistic gravity and gravitational radiation \\citep{KramerWex2009}. The system has several strange features that challenge the current understanding of pulsars and provide an uncommon opportunity to improve pulsar theories. One of the most interesting properties is the observed modulation of each pulsar's signal by the energy flux from the other, as evidenced by each pulsar's modulation period being approximately equal to the other pulsar's pulse period \\citep{McLaughlin04a,McLaughlin04b}. \\citet{Freire09a} proposed  a technique for analysing the arrival times of the pulsars' pulses and their mutual modulations which could yield the sense of rotation of each pulsar with respect to its orbital motion, among other quantities. In this paper, a complementary technique is presented with the same objective, but using measured periods rather than arrival times.   The principal benefit of our approach is that it is simpler and more intuitive.  If the validity of either of these techniques is confirmed, not only will new insights be gained but also the correctness of the lighthouse model will be empirically assessed beyond dispute. \\begin{table*} \\centering \\caption{Earth-Sun  and double pulsar  analogies. All quantities are defined as magnitudes.} \\begin{minipage}{184mm} \\begin{tabular}{lcc} \\hline Rotating object:& Earth $(\\earth)$&Pulsar A \\\\ \\hline Orbital period, Mean orbital frequency  &   $P_{\\rm orb} (\\equiv$ year), $f_{\\rm orb}$ &  $P_{\\rm orb}, f_{\\rm orb}$\\\\ \\hline \\hspace{0.4in} Instantaneous  orbital  frequency   &   \\hspace{0.7in} $f_{\\rm orb}(t)$  &  \\hspace{0.43in}  $ f_{\\rm orb}(t)  $    \\\\ \\hline Rotating object's  sidereal spin period, frequency  &   $P_{\\earth,0} (\\equiv$ sidereal day), $f_{\\earth,0}$ &   $P_{\\rm A,0}, f_{\\rm A,0} $  \\\\ \\hline Rotating object's spin period, frequency  (measured at other body) &   $P_{\\earth\\ \\rm at\\  \\sun} (\\equiv$solar day), $f_{\\earth\\ \\rm at\\  \\sun}$ &   $P_{\\rm A\\ at\\  B}, f_{\\rm A\\ at\\  B}$  \\\\  $\\equiv\\ $Modulation period, frequency at other body due to rotating object   & $\\equiv P_{\\rm m\\ at\\  \\sun\\ due\\ to\\ \\earth},$  &   $\\equiv P_{\\rm m\\ at\\  B\\ due\\ to\\ A}, $ \\\\ &\\ \\ \\ $ f_{\\rm m\\ at\\  \\sun\\ due\\ to\\ \\earth} $&  \\ \\ \\ $f_{\\rm m\\ at\\  B\\ due\\ to\\ A}$  \\\\  \\hline \\end{tabular} \\end{minipage} \\label{tab:one} \\end{table*} ",
        "conclusions": "\\cite{McLaughlin04a} presented a modulation pattern similar to drifting subpulses in the signal of PSR B, but with a frequency of 44 Hz, close to the pulse frequency of PSR A. The presence and sense of rotation of A is encoded in the observed modulation pattern, and can be revealed through an arrival-time-based analysis \\citep{Freire09a} or the frequency-based analysis presented above.  Our procedure offers the benefits of relative conceptual simplicity and close analogy with familiar phenomena in the Earth-Sun system. We present a frequency-based procedure,  building upon that used in binary pulsar search software, to distinguish among direct, retrograde, or no rotation of PSR A by searching synchronously for one of the three possible modulation signals over the full span of  available data.  Although the lighthouse model has been widely accepted, there has nevertheless been no direct observational evidence in its support up to this time. A  strength of our technique is its ability to provide such a test, empirically supporting or refuting the model.   \\citet{ferdman13} have shown that  A's spin and orbital axes are aligned to within $3 \\degr$, but they have no direct means of distinguishing parallel from antiparallel alignment.  By assuming parallel alignment, they are able to conclude that the second supernova, which created B, was relatively symmetric. Therefore the presence and sense of the rotation, as revealed by this analysis, will further test their and others' (e.g., \\citet{kramerstairs08,farret11}) evolutionary scenarios.    While \\citet{McLaughlin04b} also found that A's signal is modulated at B's frequency near A's eclipse, the mechanism is different and not dependent on B's rotation.  If, however, the modulation is observed in the future away from eclipse, it will provide an easier test for B's rotation than does the approach delineated above for A.  Our analysis can also be used on other binary systems discovered to possess phenomena caused by mutual interactions."
    },
    "1402/1402.4300_arXiv.txt": {
        "abstract": "We present a photometrical and morphological study of the properties of low redshift (z $<$ 0.5) quasars based on a large and homogeneous dataset of objects derived from the Sloan Digital Sky Survey (DR7). This study over number by a factor $\\sim$ 5 any other previous study of QSO host galaxies at low redshift undertaken either on ground or on space surveys. We used $\\sim$ 400 quasars that were imaged in the SDSS Stripe82 that is up to 2 mag deeper than standard Sloan images. For these quasars we undertake a study of the host galaxies and of their environments. In this paper we report the results for the quasar hosts.  We are able to detect the host galaxy for more than  3/4 of the whole dataset and characterise the properties of their hosts. We found that QSO hosts are dominated by luminous galaxies of absolute magnitude M*-3 $<$ M(R) $<$ M*. For the unresolved objects we computed a upper limit to the host luminosity. For each well resolved quasar we are also able to characterise the morphology of the host galaxy that turn out to be more complex than what found in previous studies. QSO are hosted in a variety of galaxies from pure ellipticals to complex/composite morphologies that combine spheroids, disk, lens and halo.  The black hole mass of the quasar, estimated from the spectral properties of the nuclei, are poorly correlated with the total luminosity of the host galaxy. However, taking into account only the bulge component we found a significant correlation between the BH mass and the bulge luminosity of the host. ",
        "introduction": "Accretion onto a supermassive black hole (SMBH) is the main mechanism that sustains the powerful activity of active galactic nuclei but may also represent a common phase in the evolution of normal galaxies. A number of fundamental question about the formation of the QSO phenomenon like the fuelling and triggering mechanisms are strictly related to the immediate environments of the active nucleus and in particular to its host galaxy \\citep{merloni10}. SMBHs may well have a period of maximum growth (maximum nuclear luminosity) contemporaneous with the bulk of the initial star formation in the bulge of galaxies. Studies of the co-evolution of SMBH and their host spheroids are therefore obviously critical to understanding how and when galaxies in the local Universe formed and evolved. The last ten years have yielded considerable progress in characterising AGN host galaxies. At variance with inactive galaxies their study is often hampered by the presence of the luminous central source that outshines the light of the host galaxy. A problem that becomes more serious for high luminosity AGNs and for sources at high redshift.  In spite of these limitations the characterization of the properties of the host galaxies offers the unique opportunity to investigate the link between the central black hole mass and its host galaxy at moderate to high redshift and to trace the possible co-evolution at different cosmic epochs. This is because for broad line AGN like quasars it is possible to estimate the mass of the central BH using kinematic arguments that are not directly dependent on the host galaxies properties.   Both ground-based and HST studies have shown that virtually all luminous low redshift (z$<$0.5) quasars reside in massive, spheroid-dominated host galaxies, whereas at lower luminosities quasars can also be found in early- type spiral hosts (e.g. Bahcall et al. 1997; Dunlop et al. 2003; Pagani et al. 2003; Floyd et al. 2004; Jahnke et al. 2004). This is in good agreement with the BH -- bulge relationship in inactive galaxies (e.g. Gultekin et al. 2009), since very massive BHs power luminous quasars. Only a small fraction of the host galaxies ($\\sim$15\\% ) are found in merger systems but it is difficult to determine clear merger signatures from morphology alone. At low redshifts a major contribution to the properties of quasar host galaxies has been provided by images from the Hubble Space Telescope (HST). The improved spatial resolution has allowed the characterization of the structure and the detailed morphology of the host galaxies \\citep{bahcall97, kukula01, ridgway01,dunlop03,peng06,zakamska06}. It turned out that QSO are hosted in luminous galaxies that are often dominated by the spheroidal component.  At high redshift (z $>$1) HST observations of quasar host galaxies (e.g. \\cite{peng06,floyd13} and references therein) have been complemented by significant contributions from 8-m class ground-based telescopes under superb seeing conditions \\citep{kotilainen07,kotilainen09} and/or with adaptive optics  \\citep{falomo08}. Comparison of host galaxies of AGN at high and low redshift constrain host galaxy evolution, as compared with the evolution of normal (inactive) galaxies.   Most of the old studies of quasar host considered few tens of objects therefore in order to derive a picture of the host properties at various redshift one should combine many different samples often obtained with different telescopes and filters. Observations carried out by HST are certainly more homogeneous (although different filters were used) and allow to investigate  a somewhat large sample based on high quality data. Nevertheless the size of these samples remain relatively small For instance in the range 0.25 $ < z < $ 0.5 about 50 QSO were imaged by HST (see references above).  In order to explore a significantly larger dataset of QSO one should refer to large surveys that include both imaging and spectroscopic data. In this respect one of the most productive recent surveys is the Sloan Digital Sky Survey that allowed to find  105783 quasars \\citep{schneider2010} from (DR-7).  Standard SDSS images are, however, too shallow and the faint nebulosity around the nucleus of quasars is not  detected. This problem has been overcome in the case of the special sky region mapped by SDSS for the SDSS Legacy Survey.  The central stripe in the South Galactic Gap, namely the Stripe82 \\citep{annis2011} is a stripe along the Celestial Equator in the Southern Galactic Cap. It is 2.5$^{\\deg}$ wide and covers -50$^{\\deg}\\leq$RA$\\leq$+ 60$^{\\deg}$, so its total area is 275deg$^2$. Stripe 82 was imaged by the SDSS multiple times, these data were taken in 2004 only under optimal seeing, sky brightness, and photometric conditions (i.e., the conditions required for imaging in the main Legacy Survey; York et al. (2000)). In 2005-2007, 219 additional imaging runs were taken on Stripe 82 as part of the SDSS supernova survey \\citep{frieman08}, designed to discover Type Ia supernovae at 0.1$< z <$0.4. The total number of images reaches $\\sim$100 for the S strip and $\\sim$ 80 for the N strip. The final frames were obtained by co-adding selected fields in r-band, with seeing (as derived from 2D gaussian fit of stars and provided by SDSS pipeline) better than 2\", sky brightness $\\leq$19.5 $mag/arcsec^2$ and less than 0.2 mag of extinction. In this area there are 12434 quasars.  Recently \\cite{matsuoka14} analyzed the stellar properties of about 800 galaxies hosting optically luminous, unobscured quasars at z$<$ 0.6 using Stripe82 images. They focused on the color of the host galaxies and found that the quasar hosts are very blue and almost absent on the red sequence with a marked different distribution from that of normal (inactive) galaxies.  For our study we selected QSO with redshift less than 0.5 for which the stripe 82 images allow us also to study the QSO galaxy environments. We adopt the concordance cosmology with H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.3 and $\\Omega_\\Lambda$ = 0.7.  In this first paper of a series we focus on the properties of quasar hosts and their relationship with the central BH mass in the explored redshift range. In forthcoming papers we investigate the galaxy environments \\citep{karhunen13}  and galaxy peculiarities \\citep{bettoni14}. A preliminary account of these results was presented in \\cite{kotilainen13}. ",
        "conclusions": "We have investigated  the properties of the host galaxies from a large ($\\sim$ 400 objects) and homogeneous dataset of low redshift (z $<$ 0.5) quasars using th SDSS images in the Stripe82 region that are significantly deeper that standard SDSS data. The 2D analysis of the images allowed us to well resolve the quasar host for  3/4 of the objects in the sample, marginally resolve other 40 quasars and derive limits for the galaxy luminosity for the unresolved targets (60 objects).  The following  properties of quasar hosts are derived:  \\begin{enumerate} \\item the luminosity of the host galaxies of low z quasars span a range from M(R) $\\sim$ -21.5 to M(R) $\\sim$ -24.0; the bulk of the host galaxies are located in the region corresponding to M*-1 and M*-2; there is a mild increase of the host luminosity with the redshift that is consistent with the passive evolution of the underlying stellar population  \\item the morphology of the host galaxies turned out to be rather complex with both bulge and disc dominated galaxies; about one third of the objects in our sample show features characteristics of bulge and disc components  \\item irrespective of the host morphology the size of the galaxies (as derived from the half light radius) ranges from compact (few kpc) objects up to extended galaxies (10-15 kpc); in the observed redshift range we do not find any significant trend  of change of the galaxy size with z  \\item  the nuclear and host galaxy luminosities are not correlated suggesting that accretion rate, BH mass, and galaxy masses and morphology combine together to smear significantly the correlation between BH and host masses  \\item the BH mass of quasars estimated from the QSO continuum luminosity and the width of the broad emission lines is poorly correlated with the total luminosity/mass of the whole host galaxy; on the contrary when the fraction of bulge to disc component is considered we find a significant correlation between the BH mass and the bulge luminosity of the host.  \\end{enumerate}  Another important source of information to characterise the properties of low redshift QSO come from the analysis of their galaxy environments as compared with those of similar galaxies with no active nuclei. These aspects will be pursued in forthcoming papers of this series \\citep{bettoni14,karhunen13}."
    },
    "1402/1402.5210.txt": {
        "abstract": "We present a new approach for estimating the 3.6 $\\mu m$ stellar mass-to-light ratio $\\Upsilon_{3.6}$ in terms of the [3.6]-[4.5] colors of old stellar populations.  Our approach avoids several of the largest sources of uncertainty in existing techniques using population synthesis models.  By focusing on mid-IR wavelengths, we gain a virtually dust extinction-free tracer of the old stars,   avoiding the need to adopt a dust model to correctly interpret optical or optical/NIR colors normally leveraged to assign the mass-to-light ratio $\\Upsilon$.  By calibrating a new relation between NIR and mid-IR colors of giant stars observed in GLIMPSE we also avoid the discrepancies in model predictions for the [3.6]-[4.5] colors of old stellar populations due to uncertainties in the molecular line opacities assumed in template spectra. We find that the [3.6]-[4.5] color, which is driven primarily by metallicity, provides a tight constraint on $\\Upsilon_{3.6}$, which varies intrinsically less than at optical wavelengths. The uncertainty on $\\Upsilon_{3.6}$ of $\\sim$0.07 dex due to unconstrained age variations marks a significant improvement on existing techniques for estimating the stellar M/L with shorter wavelength data.  A single $\\Upsilon_{3.6}$=0.6 (assuming a Chabrier IMF), independent of [3.6]-[4.5] color, is also feasible as it can be applied simultaneously to old, metal-rich and young, metal-poor populations, and still with comparable (or better) accuracy ($\\sim$0.1 dex) as alternatives. We expect our $\\Upsilon_{3.6}$ to be optimal for mapping the stellar mass distributions in S$^4$G galaxies, for which we have developed an Independent Component Analysis technique to first isolate the old stellar light at 3.6 $\\mu m$ from non-stellar emission (e.g. hot dust and the 3.3 PAH feature).  Our estimate can also be used to determine the fractional contribution of non-stellar emission to global (rest-frame) 3.6 $\\mu m$ fluxes, e.g. in WISE imaging, and establishes a reliable basis for exploring variations in the stellar IMF. ",
        "introduction": "} Accurate maps of the stellar mass distribution in nearby galaxies are essential for charting the dynamical influence of stellar structures over time.  They provide a key perspective on the evolution in the amount and distribution of baryons from high redshift to the local universe, both as the result of internal, secular processes (e.g. Elmegreen et al. 2007; Zhang \\& Buta 2007; Haan et al 2009; Foyle et al. 2010), and given `bottom up' assembly and external (heirarchical accretion) events within the context of CDM cosmology ( e.g., Balcells et al. 2003 and Courteau et al. 1996;  Florido et al. 2001; and see Navarro \\& White 1994; Somerville 2002; Governato et al. 2007; Guo \\& White 2008).\\\\ \\indent NIR bands are often though to be optimal windows on the old stars that dominate the baryonic mass in galaxies, with images at these wavelengths serving as good relative proxies for stellar mass (Elmegreen \\& Elmegreen 1984; Rix \\& Zaritsky 1995; Grosbol 1993). Yet even in these bands the light is affected by the presence of dust extinction and young stars (albeit less than in the optical), requiring the application of a M/L ratio that incorporates dependencies on metallicity, star formation history, and especially the age of the stellar population (Rix \\& Rieke 1993; Rhoads 1998). Such estimates have emerged only relatively recently, following the demonstrated correlation between optical colors and the mass-to-light ratio in stellar population synthesis models \\citep{bdJ01}.  The state-of-the art, which uses more than one color and incorporates a model of the dust, is optimal for use in 2D and has greatly reduced uncertainties \\citep{z09}.  But, at present few such maps exist.  \\indent We have begun an effort to map the stellar mass distribution in nearby galaxies in the Spitzer Survey of Stellar Structure in Nearby Galaxies (S$^4$G), an un-paralleled inventory of stellar mass and structure in $\\sim$2300 nearby galaxies ($D$$<$40 Mpc; Sheth et al. 2010).  The survey consists of deep imaging at 3.6 $\\mu m$ and 4.5 $\\mu m$ where the light traces the old stars, and where the effects of dust extinction and star formation activity (e.g. young stars, HII regions; see \\citealt{sheth} and references therein) are at a minimum compared to ground-based images in either optical or NIR wavebands.  Our method begins by using Independent Component Analysis (ICA) to isolate the old stellar light (from K and M giants) from the `contaminating' emission from several other sources: the 3.3 $\\mu m$ PAH emission feature in channel 1, accessory non-stellar continuum emission (i.e. from hot dust and PAH), and lower-M/L asymptotic giant branch (AGB) and red supergiant (RSG) stars. These three main contaminants can lead to an over-estimation of stellar mass when using uncorrected 3.6 $\\mu m$ and/or 4.5 $\\mu m$ images (\\citealt{meidtpaperI}; hereafter Paper I).  Corrections for this `contaminating' emission are especially critical for mass estimates, since this emission is not incorporated into SED libraries at the core of now-standard techniques for estimating the stellar M/L (e.g. \\citealt{bdJ01}; \\citealt{z09}).  Once contaminants are removed with ICA we obtain maps of the underlying distribution of old stellar light in S$^4$G images, with [3.6]-[4.5] colors that are consistent with those of M and K giants (\\citealt{pahre}; \\citealt{hunter}).   To generate 2D maps of the stellar mass distribution in principle thus requires applying a M/L that is appropriate for an old, dust-free population of stars.  Tight constraints on the stellar M/L at 3.6$\\mu m$ are possible using stellar population synthesis (SPS) models to link, e.g., optical and/or NIR colors to metallicity and age (or star formation history; SFH) variations in the underlying stellar population.  Paired with this information, contaminant-corrected S$^4$G images should be readily convertible into high-quality mass maps. But even in the latest generation of models there are lingering uncertainties:  1) convergence on the TP-AGB phase of stellar evolution that dominates NIR light at $\\sim$1 Gyr has yet to be reached among different models (cf. Charlot \\& Bruzual 2007, \\citealt{maraston05} and \\citealt{lancon}), 2) differences in the molecular line opacities assumed in template spectra lead to a variety of model predictions at mid-IR wavelengths (e.g. Peletier et al. 2012), and 3) realistic models of dust extinction (as well as the reflection and scattering of stellar light by dust) are difficult to implement at the pixel-by-pixel level, given the variety of dust geometries \\citep{z09}.  In this paper we explore estimating the M/L at 3.6$\\mu m$ in a way that is least susceptible to the uncertainties that plague state-of-the-art mass estimation techniques.  The key is our focus on mid-IR wavelengths where the emission is dominated by the oldest stars and dust extinction is minimal.  In $\\S$ \\ref{sec:ICA} we describe our ICA technique to isolate the old stellar light from non-stellar emission at 3.6 $\\mu m$.  Then in $\\S$ \\ref{sec:MLdescription} we demonstrate that the variation in stellar $\\Upsilon_{3.6}$ is intrinsically quite low and, in this case, we argue that the IRAC [3.6]-[4.5] colors of the old stars provide sufficient constraint on the M/L ($\\S$ \\ref{sec:withcolor}).  To demonstrate this we use SPS models, setting out first to improve their predictive power at mid-IR wavelengths.  Using a new calibration of the relation between the NIR and mid-IR colors of GLIMPSE giants (Appendix A), we link SPS models to realistic [3.6]-[4.5] colors in $\\S$ \\ref{sec:iraccolor}.  This allows us to explore the dependence of age/SFH (including bursts) and metallicity on the relation between $\\Upsilon_{3.6}$ and [3.6]-[4.5] color, which we present in $\\S$ \\ref{sec:therelation}.  Our approach for constructing 2D mass maps, namely, using information exclusively from 3.6 and 4.5 $\\mu m$ images revealed by ICA, is an approach that can be applied to the greatest number of galaxies in the S$^4$G sample without relying on ancillary optical and/or NIR data. ",
        "conclusions": "We have shown that the information contained in IRAC images at 3.6 and 4.5 $\\mu m$ is optimal for estimating stellar masses of nearby galaxies.  At 3.6 $\\mu m$ we combine an intrinsically modestly varying stellar M/L with the ability to constrain variations in the properties of the stellar population via the [3.6]-[4.5] colors of old stars.  We achieve an uncertainty as low as 0.06 dex by using the [3.6]-[4.5] color to constrain the metallicity-dependence of the stellar M/L; more modest unconstrained age variations at these long wavelengths contribute an uncertainty that is comparable to, if not below, that which can be achieved at shorter wavelengths, where the effects of dust extinction/reddening and young stars are present.  Much of the uncertainty normally introduced with the use of SPS models at long wavelengths ($\\lambda\\gtrsim$ 2.2 $\\mu m$) is reduced by using the observed colors of GLIMPSE giants to relate the NIR colors of SPS models to realistic [3.6]-[4.5] colors.  Our preferred method of assigning the stellar M/L uses the [3.6]-[4.5] color to account for the increase in the brightness of giant stars with increasing metallicity (traced by [3.6]-[4.5] color).  We define a relation between [3.6]-[4.5] color and $\\Upsilon_{3.6}$ that extends across the full range in [3.6]-[4.5] colors exhibited by galaxies across the Hubble sequence.  This relation is designed to be applied to the light from old stars, alone, i.e. when [3.6]-[4.5] colors reflect genuine variation in the properties of the old stellar population.  This makes it applicable to early-type galaxies, with little on-going star formation (and minimal dust extinction/reddening), or in late-type galaxies, where non-stellar contaminating emission is isolated from the old stellar light with ICA (Paper I).  Remarkably, the existence of an age-metallicity relation in old stellar populations may even lead to a preference for a single, color-independent $\\Upsilon_{3.6}$=0.6, good to an accuracy of 0.06 dex. We recommend adopting a more conservative 0.1 dex uncertainty on $\\Upsilon_{3.6}$=0.6 given a suspected atypical evolutionary history, or when it is not possible to correct for the presence of non-stellar emission (which leads to characteristically red [3.6]-[4.5] colors). In this case, additional uncertainty in the stellar mass estimate can arise as a result of the contribution from the non-stellar emission, which can be as large as 30\\% locally (i.e. an uncertainty comparable to that associated with the stellar M/L).  Either of the two estimates for $\\Upsilon_{3.6}$ presented here (with or without [3.6]-[4.5] color information) mark a significant improvement in our view of how the stellar mass is distributed in and among galaxies.  With greater confidence in the stellar M/L we also gain much-needed leverage on the dark matter content of galaxies, as well as the opportunity to explore variations in the IMF with global galaxy properties. We anticipate successfully implementing our preferred approach, one that yields the smallest uncertainty in stellar mass, to galaxies in S$^4$G; future work will first remove the contribution from non-stellar emission in S$^4$G 3.6 $\\mu m$ images with ICA and then use the [3.6]-[4.5] color of the old stellar light to assign $\\Upsilon_{3.6}$ (Querejeta et al. 2014).    \\\\ \\newline \\newline Thanks to Mariya Lyubenova for fruitful discussion and the entire S$^4$G team.  E.A., A.B., J.K., G.vdV., M.Q., S.M. and E.S. acknowledge financial support of the DAGAL network from the People Programme  (Marie Curie Actions) of the European Union's Seventh Framework Programme FP7/2007-2013/ under REA grant agreement number PITN-GA-2011-289313.  K.S., J.-C.M.-M., and T.K. acknowledge support from the National Radio Astronomy Observatory, which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. % \\newpage \\appendix"
    },
    "1402/1402.4668.txt": {
        "abstract": "\\noindent Infrared emission is an invaluable tool for quantifying star formation in galaxies. Because the 8 $\\mu$m polycyclic aromatic hydrocarbon (PAH) emission has been found to correlate with other well-known star formation tracers, it has widely been used as a star formation rate (SFR) tracer. There are, however, studies that challenge the accuracy and reliability of the 8 $\\mu$m emission as a SFR tracer.\\\\ \\indent Our study, part of the Herschel\\footnote{Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.} M33 Extended Survey (HERM33ES) open time key program, aims at addressing this issue by analyzing the infrared emission from the nearby spiral galaxy M33 at the high spatial scale of $\\sim$ 75 pc. Combining data from the \\it{Herschel} \\rm Space Observatory and the \\it{Spitzer} \\rm Space Telescope we find that the 8 $\\mu$m emission is better correlated with the 250 $\\mu$m emission, which traces cold interstellar gas, than with the 24 $\\mu$m emission. Furthermore, the L(8)/L(250) ratio is more tightly correlated with the 3.6 $\\mu$m emission, a tracer of evolved stellar populations and stellar mass, than with a combination of H$\\alpha$ and 24 $\\mu$m emission, a tracer of SFR. The L(8)/L(24) ratio is highly depressed in 24 $\\mu$m luminous regions, which correlate with known HII regions.\\\\ \\indent We also compare our results with the dust emission models by \\citet{draine07}. We confirm that the depression of 8 $\\mu$m PAH emission near star-forming regions is higher than what is predicted by models; this is possibly an effect of increased stellar radiation from young stars destroying the dust grains responsible for the 8 $\\mu$m emission as already suggested by other authors. We find that the majority of the 8 $\\mu$m emission is fully consistent with heating by the diffuse interstellar medium, similar to what recently determined for the dust emission in M31 by \\citet{draine13}. We also find that the fraction of 8 $\\mu$m emission associated with the diffuse interstellar radiation field ranges between $\\sim$60\\% and 80\\% and is 40\\% larger than the diffuse fraction at 24 $\\mu$m. ",
        "introduction": "\\noindent In order to understand the evolution of galaxies in the Universe, accurate measures of star formation (SF) within galactic structures need to be obtained. A large number of  tracers of star formation in a galaxy have been defined in the literature, using different regions of the electromagnetic spectrum. Among these, the infrared (IR) radiation is a classical tracer of activity in galaxies \\citep{kennicutt98, ken12, calz12}. In star-forming galaxies, young, massive stars are responsible for most of the ultraviolet (UV) radiation and, in the presence of dust a smaller or larger fraction of their light may be absorbed and reradiated in the IR regime. By observing galaxies in IR light, determinations about the rate and areas of star formation can be surmised.\\\\ \\indent The mid-IR wavelength region ($\\approx$7-40 $\\mu$m) in general, and the emission around $\\sim$8 $\\mu$m in particular, are among the favored tracers of recent star formation, because of their ready detectability in galaxies at high redshift (e.g., Daddi et al. 2005, Reddy et al. 2010, 2012). The emission in the  $\\sim$8 $\\mu$m region is mainly contributed by a combination of stellar photospheric emission, the featureless continuum of hot dust emission, and the Polycyclic Aromatic Hydrocarbon (PAH) spectral features (e.g., \\citet{smith07}). As the PAH features contribute about 70\\% or more of the emission in the Spitzer $\\sim$8 $\\mu$m band \\citep{smith07}, we call the emission in this band `PAH emission' henceforth. \\citet{roussel01} and \\citet{forster04} have suggested, using data from ISO,  that the PAH emission is closely related to the hot dust emission at 15 $\\mu$m and other tracers of star formation. Other studies (e.g. \\citet{bos04,haas02}) and follow up investigations, utilizing the \\it{Spitzer} \\rm Space Telescope \\citep{werner04}, found the relation to be more complex than previously inferred. The PAH emission does not immediately correlate with star formation as traced by the 24 $\\mu$m emission,  as the L(8)/L(24) luminosity ratio is depressed in regions of known star formation relative to the ratio in the diffuse starlight field \\citep{helou04, calz05, bendo08, pov07}. \\citet{bendo06, bendo08}, furthermore, showed that the PAH emission is more closely related to tracers of cool dust emission, heated by the diffuse starlight field, than tracers of warm dust, heated by recent star formation. This is confirmed by \\citet{verley09}, who used a complementary approach to our own to determine that at least 60\\% of the 8 and 24 $\\mu$m emission in M33 is diffuse. Conversely, a recent study by \\citet{crock13} determined that the 8 $\\mu$m emission in the galaxy NGC 628 has only a 30\\% - 43\\% fraction unassociated with recent star formation. These contradictory results call for further studies to assess the validity of the PAH emission as a SFR tracer.\\\\ \\indent The powerful \\it{Herschel} \\rm Space Observatory \\citep{pilbratt10} has targeted nearby galaxies with unprecedented resolution at the infrared wavelengths ($\\geq 70 \\mu$m) where the warm dust emission, powered primarily by massive stars, is progressively supplanted by the cold dust emission powered by low mass stars \\citep{bendo10,bendo12,boq11}. The synergy between the \\it{Spitzer} \\rm and \\it{Herschel} \\rm imaging data is such that the infrared spectral energy distribution of individual star forming regions from $\\sim$3 $\\mu$m to the sub-mm can be separated from that of the diffuse starlight in nearby galaxies, permitting perusal of the emission originating in these regions. With this level of detail, the properties of dust emission and its relation to the SFR can be better investigated. This paper uses data from both \\it{Spitzer} \\rm and \\it{Herschel} \\rm to observe the nearby galaxy, M33, in a range of IR wavelengths.\\\\ \\indent M33 is a member of the Local Group and a spiral galaxy, with an inclination of 56$^\\circ$ \\citep{regan94} and a distance of 840 kpc \\citep{freedman91}. Because of its proximity, M33 is an ideal site for the investigation of the properties of dust emission: 1$^{\\prime\\prime}$ subtends a spatial scale of $\\sim$4 pc. M33 harbors a large number of HII regions \\citep{bou74,hodge99,verley10, rel13}, which are easily identifiable in the \\it{Spitzer} \\rm  and \\it{Herschel} \\rm images. M33 has an oxygen abundance of about half solar and a shallow--to--negligible metallicity gradient as a function of galactocentric distance \\citep{ros08, magrini09,bre11}; this characteristic enables us to investigate the PAH emission with less attention to the effect of metallicity on the strength of the PAH features \\citep{hunt05, eng05, mad06,draine072, eng08}. The \\it{Herschel} \\rm data for this project came from the HERM33ES open time key program \\citep{kramer10} and \\it{Spitzer} \\rm data were obtained through the Guaranteed Time observations for the IRAC and MIPS instruments \\citep{hinz04, mcquinn07}. We also make use of the H$\\alpha$ image from \\citet{hoopes00}, taken with the 0.6 meter Burrell-Schmidt telescope at Kitt Peak National Observatory. Details about the imaging and its reduction can be found in that paper.\\\\ \\indent Relationships among the PAH 8 $\\mu$m, 24 $\\mu$m, 250 $\\mu$m, and total infrared (TIR) emission are the primary subjects of this study, under the reasonable assumption that the 24 $\\mu$m emission is mainly tracing current SF \\citep{calz05,calz07} and that the 250 $\\mu$m emission is mainly tracing cold dust.\\\\ \\indent Our analysis plan is similar to that used by \\citet{bendo08} for the SINGS galaxies \\citep{ken03}: we will investigate the relationship of the  8 $\\mu$m emission with tracers of the warm  (24 $\\mu$m) and cool (250 $\\mu$m) dust emission, and with the total infrared (TIR) emission. The advantage of our analysis over that of Bendo et al. is twofold: (1) M33 is 3 to 20 times closer than the SINGS galaxies, thus enabling exquisite spatial resolution; and (2) the 250 $\\mu$m emission from {\\it Herschel} traces more closely the cool dust and has more than twice the spatial resolution of the 160 $\\mu$m {\\it Spitzer} data used by Bendo et al.\\\\ \\indent In what follows, L($\\lambda$) refers to monochromatic luminosity derived as \\begin{equation} L(\\lambda) = [\\nu L_{\\nu}]_{\\lambda} \\end{equation} in units of erg s$^{-1}$ kpc$^{-2}$. In this work, each pixel will be $\\sim$73.3 pc in size, or about the size of a large HII region. ",
        "conclusions": "\\noindent In this paper, we have investigated the relationships between the 8 $\\mu$m emission and emission in other wavelengths that are correlated to old stellar populations, star formation, and cold dust. We have also compared our data with the predictions of the models from \\citet{draine07}. Ratios with 24 $\\mu$m and TIR luminosity show the 8 $\\mu$m emission to become underluminous in areas of strong 24 $\\mu$m emission. Taking the 24 $\\mu$m emission to be the product of heating of dust by young stellar radiation, this implies that 8 $\\mu$m emission is not a close tracer of young stellar populations, and the carriers responsible for its emission are possibly being destroyed by the intense radiation fields of star forming regions. Furthermore, the behavior of most ($\\ge$ 80\\%) 8 $\\mu$m emission can be explained using only the U$_{min}$ = U$_{max}$ models, which is valid for diffuse ISRF. We also find that the values of U$_{min}$ that account for the observed 8~$\\mu$m emission are in the range 0.1 - 25, consistent with the range found by Draine et al. (2014) for the galaxy M31. When applying a model that includes heating from HII regions, the observed 8 $\\mu$m values are lower than expected, supporting other authors' conclusions that the carriers responsible for 8 $\\mu$m emission are destroyed in intense stellar radiation fields. An alternative explanation to the destruction of the 8 $\\mu$m carriers is to make the PAHs more neutral, suppressing the 8 $\\mu$m emission and shifting power to the 11-12 $\\mu$m features\\citep{sand12}. Mid-IR spectroscopy may discriminate among these two scenarios, although we tentatively give preference to the PAH destruction interpretation.\\\\ \\indent Ratios of 8 $\\mu$m with 250 $\\mu$m emission show a strong correlation with both the TIR and 3.6 $\\mu$m luminosities. Both 3.6 $\\mu$m emission and L(8)/L(250) show a similar galactocentric radial trend, suggesting the 8 $\\mu$m emission may be more connected to old stellar populations than star formation. Further, the L(8)/L(250) as a function of 3.6 $\\mu$m luminosity shows the tightest relation among those presented in this paper, especially at high luminosity.\\\\ \\indent We also derive the fraction of 8 $\\mu$m and 24 $\\mu$m emission heated either by the ISRF or HII regions. The fraction of L(24) associated with diffuse emission is 60 - 43\\%, while the fraction of L(8) is nearly 33 - 42\\% higher, with 80 - 59\\% of the emission coming from heating by the ISRF.\\\\ \\indent The results of this study, which support findings by other authors, but at the exquisite spatial resolution enabled by the proximity of M33, suggest that the 8 $\\mu$m luminosity should not be used as a proxy for measuring and locating star formation in galaxies. Emission in this wavelength is shown in M33 to be more correlated with cold dust or old stellar populations than with star formation."
    },
    "1402/1402.4136_arXiv.txt": {
        "abstract": "{We present a study of the overdensity of X-ray-selected active galactic nuclei (AGN) in 33 galaxy clusters in the XMM-LSS field (The XMM-Newton Large Scale Structure Survey), up to redshift $z=1.05$ and further divided into a lower ($0.14\\leq z\\leq 0.35$) and a higher redshift ($0.43\\leq z\\leq1.05$) subsample. Previous studies have shown that the presence of X-ray-selected AGN in rich galaxy clusters is suppressed, since their number is significantly lower than what is expected from the high galaxy overdensities in the area. In the current study we have investigated the occurrence of X-ray-selected AGN in low ($\\langle L_x, bol\\rangle=2.7\\times10^{43}$ erg/s) and moderate ($\\langle L_x, bol\\rangle=2.4\\times10^{44}$ erg/s) X-ray luminosity galaxy clusters in an attempt to trace back the relation between high-density environments and nuclear activity. Owing to the wide contiguous XMM-LSS survey area, we were able to extend the study to the cluster outskirts. We therefore determined the projected overdensity of X-ray point-like sources around each cluster out to $6r_{500}$ radius, within $\\delta r_{500}=1$ annulus, with respect to the field expectations based on the X-ray source $\\log N -\\log S$ of the XMM-LSS field. To provide robust statistical results we also conducted a consistent stacking analysis separately for the two $z$ ranges. We investigated whether the observed X-ray overdensities are to be expected thanks to the obvious enhancement of galaxy numbers in the cluster environment by also estimating the corresponding optical galaxy overdensities, and we assessed the possible enhancement or suppression of AGN activity in clusters.  We find a positive X-ray projected overdensity in both redshift ranges at the first radial bin, which however has the same amplitude as that of optical galaxies. Therefore, no suppression (or enhancement) of X-ray AGN activity with respect to the field is found, in sharp contrast to previous results based on rich galaxy clusters, implying that the mechanisms responsible for the suppression are not as effective in lower density environments. After a drop to roughly the background level between 2 and $3 r_{500}$, the X-ray overdensity exhibits a rise at larger radii, significantly greater than the corresponding optical overdensity. The radial distance of this overdensity \\textquotedblleft bump\", corresponding to $\\sim 1.5 - 3$ Mpc, depends on the richness of the clusters, as well as on the overall X-ray overdensity profile.  Finally, using the redshift information, photometric or spectroscopic, of the optical counterparts, we derive the spatial overdensity profile of the clusters. We find that the agreement between X-ray and optical overdensities in the first radial bin is also suggested in the 3-dimensional analysis. However, we argue that the X-ray overdensity \\textquotedblleft bump\" at larger radial distance is at least partially a result of flux boosting by gravitational lensing of background QSOs, confirming previous results. For high-redshift clusters the enhancement of X-ray AGN activity in their outskirts appears to be intrinsic. We argue that a spatial analysis is crucial for disentangling irrelevant phenomena affecting the projected analysis, but we are still not able to report statistically significant results on the spatial overdensity of AGN in clusters or their outskirts because we lack the necessary numbers.} ",
        "introduction": "As one of the most powerful extragalactic phenomena, active galactic nuclei (AGN) are a valuable tool in the study of the universe, since they can be used as cosmological probes, provide answers to various problems of galaxy evolution, and shed light on the innermost regions of galaxies, where the super massive black hole resides and highly energetic processes take place. However, the triggering mechanism of the omnipresent, but not always active, black hole is still elusive. Although major merging of gas-rich galaxies seems to be a highly probable mechanism for the triggering of nuclear activity (e.g., Sanders et al. 1988; Barnes \\& Hernquist 1991; Hopkins et al. 2006), recent studies of the morphology of AGN hosts find no evidence of a major merging-AGN connection (e.g., Cisternas et al. 2011; Kocevski et al. 2012). On the other hand, minor merging and interactions are still strongly disputed, while secular evolution also seems able to feed the central engine since many AGN are found to be isolated and undisturbed (e.g., Hopkins \\& Hernquist 2006; Cisternas et al. 2011; Kocevski et al. 2012). Therefore, the effect of the environment on the activity of the nucleus and vice versa is still fairly undetermined, but nevertheless crucial. Galaxy clusters represent the one end of the density spectrum in our universe, and as such it is an ideal place to investigate the effects of dense environment in the triggering of AGN, especially since an excessive number of X-ray point-like sources are undoubtedly found there (e.g., Cappi et al. 2001; Molnar et al. 2002; Johnson et al. 2003; D'Elia et al. 2004; Gilmour et al. 2009). Specifically, for the XMM-LSS field investigated in the current study, Melnyk et al. (2013) have found that 60\\% of X-ray-selected AGN reside in the overdense regions of group-like environment.  Theoretically the feeding of the black hole can only be achieved by means of a non-axisymmetric perturbation that induces mass inflow. This kind of perturbation can be provided by interactions and merging between two galaxies, and the result of the inflow is the feeding of the black hole and activation of the AGN phase (e.g., Umemura 1998; Kawakatu et al. 2006; Koulouridis et al. 2006a, 2006b, 2013, Ellison et al. 2011; Silverman et al. 2011; Villforth et al. 2012; Hopkins \\& Quataert 2011). Thus, the cluster environment, where the concentration of galaxies is very high relative to the field, would also seem favorable to AGN. However, the rather extreme conditions within the gravitational potential of a galaxy cluster can work in the opposite direction as well. Ram pressure from the inter cluster medium (henceforth ICM) able to strip/evaporate the cold gas reservoir of galaxies (Gunn \\& Gott 1972; Cowie \\& Songaila 1977; Giovanelli et al. 1985) can strongly affect the feeding of the AGN. Nevertheless, other studies have argued that ram pressure stripping cannot be as effective in transforming blue sequence galaxies to red (e.g., Larson et al. 1980; Balogh et al. 2000, 2002; Bekki et al. 2002; van den Bosch et al. 2008; Wetzel et al. 2012), especially in lower density clusters where other processes should take place as well. Large velocity dispersion of galaxies within clusters could also prevent the effective interactions (Aarseth \\& Fall 1980), particularly mergers, while the fast \\textquotedblleft grazing\" bypassing galaxies may also cause gas stripping by \\textquotedblleft harassment\" (e.g., Natarajan et al. 2002, Cypriano et al. 2006). However, the efficiency of this phenomenon has once more been questioned (e.g., Giovanardi et al. 1983). A combination of the above mechanisms, in addition to the possible prevention of accretion of halo mass into cluster galaxies (\\textquotedblleft strangulation\"; e.g., Larson et al. 1980; Bekki et al. 2002; Tanaka et al. 2004) may, in fact, suppress the AGN activity in clusters despite the number of potentially merging and interactive galaxies.  When using only optical data, the results seem to remain inconclusive. Early studies reported that AGN are less frequent in galaxy clusters than in the field (Osterbrock 1960; Gisler 1978; Dressler, Thompson \\& Schectman 1985) and recent large-area surveys support this suggestion (Kauffmann et al. 2004; Popesso \\& Biviano 2006; von der Linden et al. 2010; Pimbblet et al. 2013). Other studies, however, have found no differences between clusters and field (e.g., Miller et al. 2003), at least when selecting the weak AGN (e.g., Martini et al. 2002; Best et al. 2005; Martini et al. 2006; Haggard et al. 2010). We should note here that considering only the optical wavelengths is not the optimal way of finding AGN since they suffer greatly from absorption. Especially if gas depletion is at play and low accretion rates are expected, then  most of the spectral signatures of the AGN could be \\textquotedblleft buried\" in the host galaxy.  Radio-loud AGN on the other hand, seem to be more clustered than any other type of galaxy (Hart, Stocke \\& Hallman 2009) and are often associated with BCGs (brightest cluster galaxies)(e.g., Best 2004; Best et al. 2007). In addition, the fraction of X-ray AGN in BCGs is higher than in other cluster galaxies (e.g., Hlavacek-Larrondo et al. 2013, and references therein). These findings can be attributed to hot gas accretion from the hot X-ray cluster halo, although gas from any other source fueling the black hole at low accretion rates would also have the same effect. If the hot gas accretion is a possible fueling mechanism for the X-ray AGN, as well, then we should expect them to reside primarily within clusters.  Undoubtedly, the best way to detect active galaxies is through X-ray observations (e.g., Brandt \\& Alexander 2010). However, during the previous decade, only a small fraction of X-ray point-like sources in clusters had positive confirmation as true cluster members (see Martini et al. 2002; Davis et al. 2003; Finoguenov et al. 2004; Arnold et al. 2009), leaving the question of whether the positive X-ray overdensities found in galaxy clusters represent enhancement or suppression of the nuclear activity unanswered. More recent studies, however, report more conclusive results by comparing X-ray to optical data. Koulouridis \\& Plionis (2010) demonstrate the significant suppression of X-ray-selected AGN in 16 rich Abell clusters (Abell et al. 1958) by comparing the X-ray point source overdensity to the optical galaxy overdensity. Ehlert et al. (2013; 2014) argue that the X-ray AGN fraction in the central regions of 42 of the most massive clusters known to date is about three times lower than the field value using the same technique. More importantly, after having complete spectroscopy for their X-ray point source sample, Haines et al. (2012) argue that X-ray AGN found in massive clusters are an infalling population, which is \\textquotedblleft extinguished\" later, and confirm the suppression in rich clusters. On the other hand, Martini et al. (2013) argue that this trend is not confirmed for a sample of high-redshift clusters ($1.0<z<1.5$), where the presence of luminous X-ray AGN ($L_{(0.5-2{\\rm\\; keV})}>10^{43}$ erg/s) is consistent with the field. We note, however, that the high-redshift regime studied and the large AGN photometric redshift uncertainties ($\\sigma_z=0.12(1+z)$, double that of normal galaxies) introduce some level of uncertainty to the results. Neverteless, they agree with findings from the DEEP2 Redshift Survey\\footnote{http://deep.ps.uci.edu/} that show that only below $z$=1.3 does the fraction of blue galaxies in groups drop rapidly and become constant below $z$=1 (Gerke et al. 2007), while the red fraction correlates weakly with overdensity above $z$=1.3 (Cooper et al. 2007). In the present study we only deal with clusters $z<$1.05, where the cluster's population is dominated by early type red galaxies. Finally, we should also mention that an indirect way to address the issue is by X-ray clustering analyses, but still their results also remain inconclusive (see relevant discussion in Haines et al. 2012 \\S5.2).   Considering the above, there is still the need to clarify the influence of the environment on the AGN phenomenon. And while the majority of the above studies are dealing with the most massive and rich clusters, the population of moderate-to-poor clusters is still overlooked.\\footnote{We should note that the categorization of galaxy clusters to different richness classes is not explicit, and it is safer to be used statistically. Nevertheless, the relation between mass and X-ray properties is well studied (e.g., Edge \\& Stewart 1991a, b; Finoguenov, Reiprich \\& Bohringer, 2001) and, also considering more recent studies of X-ray luminous clusters (e.g., Ebeling et al. 2010), we can infer that massive clusters have X-ray luminosities higher than $\\sim5\\times10^{44}$ erg/s and temperatures higher than $kT>5$ keV.} If the reason for the deficiency of X-ray AGN in rich clusters is the strong gravitational potential, which provides the necessary conditions for the suppression (whichever these may be: gas stripping, strangulation, tidal stripping, evaporation, high velocity dispersion, etc.), one would expect the AGN presence to rise in shallower gravitational potentials. A similar relation between the strength of the gravitational potential and star formation quenching (Popesso et al. 2012) supports the above expectation (see also Wetzel et al. 2012). Another issue is the radial extent of the search for X-ray AGN around clusters. An enhancement of AGN activity is observed far from the cluster's center (e.g., Fassbender et al. 2012), and it could be due to an infalling population (Haines et al. 2012) coming from the \\textquotedblleft outskirts\" of the clusters where the concentration of galaxies is still high. The question is where should we place the \\textquotedblleft outskirts\" and to what extent. Most studies could not reach farther than a 2$r_{500}$ radius, although the overdensity profile of optical galaxies remains higher than the field level even beyond that radius (e.g., Ehlert et al. 2013). Finally, what is also overlooked is the background overdensity of X-ray sources in the area of clusters. In Koulouridis et al. (2010), we used SDSS (Sloan Digital Sky Survey) optical data for all the detected X-ray point-like sources within a 1 Mpc radius, and argue that their positive overdensity values were associated with background QSOs rather than cluster members. A possible cause is the gravitational lensing of background sources, which is unimportant when compared to the large number of optical galaxies in clusters but can become very important for X-ray sources and affect the assessment of their clustering.  In the current study, our aim is to investigate the AGN phenomenon in the environment of moderate and poor clusters located in the XMM-LSS contiguous field of 11.1 deg$^2$. We identify all possible X-ray AGN candidates, which we define as sources with $L_{(0.5-2{\\rm\\; keV})}>10^{42}$ erg/s at the redshift of the cluster, and compare their overdensity in the area of the clusters to the respective overdensity of optical galaxies, available by the CFHT legacy survey. Such a large contiguous area gives us the unique opportunity not only to use a large cluster sample but also to extend our search for X-ray AGNs around clusters at great distances, reaching homogeneously up to a 6$r_{500}$ radius. In addition, we make use of photometric redshift data calculated specifically for X-ray-selected AGN hosts, in an attempt to assess the true number of X-ray AGNs in our clusters and clarify the effect of the excessive overdensity of background X-ray sources (e.g., Koulouridis \\& Plionis 2010).  We describe our samples and methodology in \\S2, while our results and conclusions are presented in \\S3 and \\S4, respectively. Throughout this paper we use $H_0=72$ km/s/Mpc, $\\Omega_m=0.27$, and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "We conducted a statistical study of 33 clusters of poor and moderate richness, within the XMM-LSS field that covers $\\sim$20$\\%$ of the XXL survey, by comparing the density of X-ray sources within multiples of the $r_{500}$ radius with the expected field density, calculated from the $\\log N - \\log S$ for the same area. We compared this projected overdensity with the respective optical galaxy overdensity in an attempt to estimate the suppression or the enhancement of X-ray-selected AGN. In addition, we calculated the spatial overdensities, using the available spectroscopic and photometric redshifts in an attempt to identify and quantify the true cluster members and explain the results of our previous projected analysis.  The conclusions that can be drawn from the above analysis are the following:  $ \\bullet$ The projected analysis of X-ray versus optical overdensity within the two central $r_{500}$ annuli, corresponding to $\\sim 1$ Mpc radius, results in a strong positive signal showing that the environment of the low and moderate X-ray-luminosity clusters of our samples does not suppress the X-ray AGN activity. This result is in sharp contrast to the outcome of many studies of rich clusters (e.g., Koulouridis \\& Plionis 2010; Ehlert et al. 2013; Haines et al. 2012), which implies that lower richness cluster environments do not suppress X-ray AGN activity. Interestingly, in even lower density environments (galaxy groups), an enhancement of X-ray AGN may be present (Melnyk et al. 2013).  $ \\bullet $ After calculating the projected overdensities at large radial distances from the center of the cluster ($3^{rd}-5^{th}$ bins, corresponding to $\\sim 1.5$-3 Mpc, depending on the redshift), a significant rise in the X-ray source overdensity is observed. This excess has also been reported in previous studies (Haines et al.2012, Fassbender et al. 2012) and has been attributed to an infalling population of galaxies from the outskirts of the clusters that interact and merge, producing the observed overdensity of X-ray AGN. This surplus is confirmed for both our low- and high-redshift clusters.  $ \\bullet $  Using spectroscopic and photometric redshifts, we discovered that the X-ray \\textquotedblleft bump\" at a large radial distance vanishes completely from the poor low-$z$ sample, and we argue therefore that this density excess may be produced by flux boosting of background sources due to gravitational lensing, sometimes even enhanced by additional background galaxy clusters along the same line of sight. On the other hand, a high X-ray source overdensity persists in the last annulus of the moderate X-ray luminosity high-$z$ sample, implying that for intermediate \\textquotedblleft -richness\" clusters, additional triggering of X-ray AGN in the outskirts is still possible.  In a nutshell, the projected overdensity analysis produces statistically significant results, but at the same time these results are contaminated by projection effects of background-lensed QSOs. On the other hand, although the spatial analysis performed is free of these effects, it is not able to reach definite results owing to the small numbers involved, making it necessary to study larger samples of galaxy clusters. Especially the area included in the annuli closer to the cluster center is so small that does not allow us to reach any definitive conclusions about the suppression of X-ray AGN. The stacking of clusters proves very useful, but splitting the total sample into two redshift subsamples again reduces the numbers greatly. However, the division is crucial since not only do we select a population of more X-ray luminous clusters in higher redshifts, but we can also detect only higher luminosity X-ray AGN.   We should stress that the large contiguous area of the XMM-LSS has allowed us to study the overdensity of X-ray AGN within large radial distances from the cluster center for the first time. This proved to be essential for exploring the relation between the dense environment of clusters and the  X-ray AGN activity in detail. To fully understand this relation, we need to trace its evolution as a galaxy approaches the cluster's gravitational potential, enters the hot ICM, and crosses the cluster. At the same time, we need to disentangle irrelevant effects such as the gravitational lensing of background sources, probably enhanced by the presence of additional clusters along the line of sight. A photometric variability study of these sources may also shed some light on whether the lensing amplification could be due to micro-lensing and/or convergence by matter in the clusters.  We believe that the analysis of the full XXL field, which is almost five times larger than the XMM-LSS (reaching 50 deg$^2$), together with a detailed spectroscopic follow-up of the optical counterparts of all X-ray point sources, detected in the XXL clusters, may provide reliable and robust results as to the origin (true enchancement, lensing, presence of background clusters, etc) of the excess X-ray sources detected in the outer $\\sim 3-5 r_{500}$ annuli of either low and high-redshift clusters."
    },
    "1402/1402.0073_arXiv.txt": {
        "abstract": "Variations in stellar flux can potentially overwhelm the photometric signal of a transiting planet. Such variability has not previously been well-characterized in the ultraviolet lines used to probe the inflated atmospheres surrounding hot Jupiters. Therefore, we surveyed 38 F-M stars for intensity variations in four narrow spectroscopic bands: two enclosing strong lines from species known to inhabit hot Jupiter atmospheres, \\ion{C}{2} $\\lm\\lm$1334,1335 and \\ion{Si}{3} $\\lm$1206; one enclosing \\ion{Si}{4} $\\lm\\lm$1393,1402; and 36.5 \\AA\\ of interspersed continuum. For each star/band combination, we generated 60~s cadence lightcurves from archival {\\it HST} COS and STIS time-tagged photon data. Within these lightcurves, we characterized flares and stochastic fluctuations as separate forms of variability. {\\bf Flares:} We used a cross-correlation approach to detect 116 flares. These events occur in the time-series an average of once per 2.5 h, over 50\\% last 4 min or less, and most produce the strongest response in \\siv. If the flare occurred during a transit measurement integrated for 60 min, 90/116 would destroy the signal of an Earth, 27/116 Neptune, and 7/116 Jupiter, with the upward bias in flux ranging from 1-109\\% of quiescent levels. {\\bf Fluctuations:} Photon noise and underlying stellar fluctuations produce scatter in the quiescent data. We model the stellar fluctuations as Gaussian white noise with standard deviation $\\s_x$. Maximum likelihood values of $\\s_x$ range from 1-41\\% for 60~s measurements. These values suggest that many cool stars will only permit a transit detection to high confidence in ultraviolet resonance lines if the radius of the occulting disk is $\\gtrsim$ 1 R$_J$. However, for some M dwarfs this limit can be as low as several $R_\\earth$. ",
        "introduction": "\\label{sec:intro}  Transit observations in the far-ultraviolet (FUV, $1200\\leq\\lambda\\leq1400$ \\AA) have revealed the existence of inflated atmospheres surrounding the hot Jupiters HD209458b \\citep{madjar03,madjar04,linsky10} and HD189733b \\citep{lecavelier10,bourrier13,jaffel13}. During the transit of these planets, the UV resonance transitions of several species in their atmospheres -- \\ion{H}{1}, \\cii, \\ion{O}{1}, and \\siii\\ -- produce detectable absorption against the background stellar emission. The depth of the absorption indicates that these species occupy a volume overfilling the planets' Roche lobes, suggesting atmospheric escape. Furthermore, such UV transit spectrophotometry can constrain the atmospheric mass loss rate (e.g. \\citealt{madjar04,lecavelier10}) and characterize the atmospheric response to changes in the stellar radiation and particle flux \\citep{lecavelier12,bourrier13}.  However, variability in the background flux source itself --  the host star -- presents additional, instrument-independent challenges for all transit observations. In a lightcurve, the transit signal can be spuriously weakened, sometimes completely obliterated, by flares buried within it, or amplified by flares flanking it. Outside of flares, the transit signal can be obscured by stochastic fluctuations in stellar luminosity that act as additional noise, compounding that of photon statistics and instrumental sources. These forms of stellar variability fundamentally limit transit observations. Therefore, evaluating the possibilities for future transit work necessitates measurements of this variability for potential host stars.  The list of potential targets is diverse, encompassing stars beyond those just on the main sequence, such as HD209458b and HD189733b. Such candidate targets could include stars with transitional or debris disks (encompassing weak-line T-Tauri stars, WTTS) where planets might be in the process of coalescing from disk material. For example, \\citet{eyken12} report on the transit signature of a super-Jupiter orbiting a WTTS star in the Orion-OB 1a/25-Ori region. High contrast imaging of stars with disks has also revealed (proto)planetary objects or evidence for these objects through disk gaps, such as LkCa15 \\citep{kraus12}, PDS 70 \\citep{hashimoto12}, RX J1633.9-2442 \\citep{cieza12}, and TW Hya \\citep{debes13}. Candidate targets also include post main-sequence stars. The evolved F5-F7 stars WASP-76, WASP-82, and WASP-90 host transiting hot Jupiters \\citep{west13}, as does the G8III giant HIP 63242 \\citep{jones13}. In addition, observations of remnant debris orbiting white dwarfs (e.g. \\citealt{farihi13}) hint at the possibility of detecting extant or decomposing planets around stars nearing the end of their lives.  To begin characterizing the range of background fluctuations faced by FUV transit observations, we have conducted a survey of stellar flares and stochastic fluctuations for the largest possible sample of stars with archival FUV photon event data from {\\it HST} covering \\cii\\ $\\lm\\lm$1334,1335 and \\siii\\ $\\lambda$1206, two key lines for probing outer atmospheres of close-in exoplanets. We also include \\siv\\ $\\lambda\\lambda$1394,1403 and a composite of interspersed continuum bands. We did not attempt to include \\ion{H}{1} $\\lambda$1216 (Ly$\\alpha$) or \\ion{O}{1} $\\lambda$1302 due to the correction for geocoronal emission that is required. This survey relies on UV data acquired by the two powerful UV spectrographs on board the {\\it Hubble Space Telescope (HST)}: the Cosmic Origins Spectrograph (COS), with its G130M and G140L gratings, and the Space Telescope Imaging Spectrograph (STIS) with its E140M grating. To enable an analysis of temporal variability, we traded spectral resolution for temporal resolution. Thus, we employed short, 60~s time bins and summed photon counts over roughly the full width ($\\sim$ 1-2 \\AA) of each line and 36.5 \\AA\\ of interspersed continuum to create lightcurves. In comparison, transit observations have achieved resolutions of  $\\sim0.1$ \\AA\\ (roughly 20-30 km s$^{-1}$) when coadding an entire transit dataset. However, these same observations commonly integrate over the full line-widths for increased signal to noise (e.g. \\citealt{madjar03,lecavelier10}).  We present the results of an analysis of 153 lightcurves of 60~s cadence, including (1) 116 flares we identified by cross-correlating with a flare-like kernel and (2) estimates or upper limits for the standard deviation of stochastic fluctuations in the quiescent portions of the lightcurve. In the remainder of this introduction, we expand upon the implications and physical sources of stellar variability and provide pointers to previous variability surveys. In Section~\\ref{sec:stars}, we describe the stellar sample and process of generating lightcurves from the FUV time-tagged photon data. In Section~\\ref{sec:analysis} we outline the variability analysis, treating flares and stochastic fluctuations separately. In Section~\\ref{sec:results}, we present the results, followed in Section~\\ref{sec:discussion} by a discussion, including the implications for observing transits. We summarize the work in Section~\\ref{sec:summary}.  Throughout this paper, we will treat 3.5~h as the typical transit timescale. This is near the average transit duration of 3.6 h (sample standard deviation 1.9 h) for exoplanets listed in the Exoplanet Data Explorer \\citep{wright11} as of 2013 October.  \\subsection{A Brief Discussion of Stellar Variability} \\label{sec:varintro}  Stellar variability can be divided into three categories based on the differing implications for transit observations: \\begin{enumerate} \\item Periodicities: Oscillations of the stellar flux by phenomena that can be sufficiently characterized to predict the level of modulation over the course of a transit. \\item Flares: Bursts of brightness isolated in time with respect to the cadence of the data, and well above the quiescent scatter in the lightcurve. \\item Stochastic Fluctuations: Variations in the stellar flux that are too chaotic to be accurately predicted on transit timescales. \\end{enumerate}  Predictable periodic signals are surmountable obstacles. Any periodic signal strong enough to be even modestly detected would allow an accurate model fit, bootstrapped over many cycles, such that removing or accounting for the signal would not obscure an otherwise detectable transit.  Strong flares complicate the lightcurve analysis. While it is possible to describe the distribution of solar flares in both strength and frequency \\citep{cassak08}, a statistical model cannot predict the onset, magnitude, and evolution of flares for any specific timeline, such as during a transit. Strong flares thus pose the risk of overwhelming, and weak flares of attenuating, a transit signal. Flares occurring near a transit could augment estimates of the out-of-transit flux, spuriously deepening the transit signal. Beyond interfering with transit observations, flares might also impact the atmospheres of exoplanets, such as the stripping of atomic hydrogen from HD189733b conjectured by \\citet{bourrier13} to be caused by an observed host-star flare.  Stochastic fluctuations, however, represent the greatest barrier to transit photometry. Because these fluctuations cannot be predicted deterministically over the course of a transit, they must be treated as noise. As with photon noise, they pose the risk of obscuring a true signal or causing a false one. This ``noise\" can be overcome by averaging measurements until the uncertainty is within tolerable limits, but with a caveat: Unlike photon noise, stochastic fluctuations are probably not white noise. For example, in 60~s cadence, broadband optical photometry from {\\it Kepler}, stochastic fluctuations of stellar flux (attributed to granulation and magnetic activity) have a power spectrum that, unlike white noise, is not flat \\citep{gilliland10}. Although {\\it Kepler} measures broadband optical flux, not the chromosphere and transition region FUV emission line flux we analyzed, the fluctuations of the chromospheric near UV flux from YZ CMi also show a frequency-dependent power spectrum \\citep{robinson99}. Therefore, it seems probable that stochastic fluctuations in stellar flux will not behave as white noise in any band. What appears as true noise at one cadence would resolve into smooth variations at some faster cadence. Below this threshold cadence, lightcurve points will be highly correlated. Binning adjacent flux measurements in this regime will not average out the scatter.  The emission line flux we analyzed samples regions with temperatures of $10^4-10^5$ K in the outer atmospheres of stars. Specifically, the peak formation temperatures of the lines in solar conditions are estimated by \\citet{dere09} to be $3.2\\sn4$ K for \\cii\\ $\\lm\\lm$1334,1335, $5\\sn4$ K for \\siii\\ $\\lambda$1206, and $8\\sn4$ K for \\siv\\ $\\lambda\\lambda$1393,1402. Reference models of the solar atmosphere place emission from these lines in the thin transition region between the chromosphere and the corona. In contrast to the lines, the solar FUV continuum (shortward of 1500 \\AA) forms at slightly lower altitudes, primarily in a region of initial chromospheric temperature rise above the photospheric temperature minimum \\citep{linsky12a}.  Variability in these regions of stellar atmospheres can result from several phenomena. Although drawing conclusions about these underlying physical phenomena is not our objective, the generally accepted origins of periodicities, flares, and stochastic fluctuations bear mentioning.  Periodic variability can be the result of a pulsational instability in the star \\citep{gautschy95, gautschy96} and/or rotation of long-lived, localized brightness variations (starspots, faculae, etc.)  through the observer's field of view \\citep{vaughan81}. A periodic signal can result from extrinsic phenomena as well. Gradual oscillations in flux are produced by phase changes of an orbiting planet (e.g. \\citealt{borucki09}). Isolated, but nonetheless periodic, dips in flux occur when an orbiting object, such as an exoplanet or stellar companion, transits the host star (e.g. \\citealt{wilson71}).  Flares are generally thought to be the result of magnetic reconnection events in the corona that abruptly convert magnetic energy into plasma kinetic energy. Some of this energy is deposited in the chromosphere and photosphere and radiated away \\citep{haisch91,gershberg05}. Flares commonly produce a sharp rise in flux followed by an exponential decay lasting from hours to minutes, possibly even seconds \\citep{pettersen89,gershberg05}. Within a single flare, multiple peaks and changes in the decay rate are possible. Some researchers have identified as flares events in which the stellar flux rises and fades more gradually \\citep{houdebine03,tovmassian03}.  The strength and frequency of flares typically exhibit an inverse power-law relationship (e.g. \\citealt{shakhovskaia89,audard00} for stellar flares, \\citealt{lin84,nita02} for solar flares). This implies that weaker flares are more prevalent than conspicuous events, such that many flares will occur in observations that cannot be clearly resolved as such. In fact, if the power law is steep enough, the lowest energy flares, often termed microflares, might inject enough heat into the corona of a star to explain the high temperatures present there \\citep{hudson91,audard00}.  Low energy ``microflares'' or even ``regular'' flares, if the data is not of sufficient quality to resolve them, will contribute to the observed stochastic fluctuations of a target. For instance, \\citet{robinson99} suggest microflaring as an explanation for the quiescent stochastic fluctuations in near UV flux that they observed from YZ CMi . They simulate the production of such stochastic fluctuations with a microflare model and find that it closely resembles the YZ CMi quiescent data. Ultimately, the extent of the contribution of flares to the observed stochastic fluctuations of any target is determined by the level of stellar magnetic activity, the photometric quality of the data, and the threshold set for identifying a lightcurve anomaly as a flare rather than fluctuation.  The remaining proportion of stochastic fluctuations in transition region line emission could be explained by several phenomena. Transition region explosions, smaller events possibly associated with magnetic restructuring at the edges of newly emerging flux loops \\citep{gershberg05}, could introduce variability while also serving as a dominant heating source for the transition region. \\citet{wood97} suggested such events might explain broad components of \\siv\\ and \\ion{C}{4} emission in the FUV spectra of 11 late-type stars. However, \\citet{peter06} suggests magnetic flux braiding and consequent Joule dissipation might be the dominant heat source for the transition region. Both braiding of surface field and the emergence of field loops produce pockets of rapid heating in the three-dimensional MHD models of \\citet{hansteen10} that could explain much of the temporal variability of line emission originating in the transition region. In the \\citet{hansteen10} model, the injected energy results from work done on the magnetic field by photospheric motions, tying transition region variability to the convective cells and p-mode oscillations within the star. These convective cells and p-mode oscillations also affect the transition region environment by initiating high altitude shock waves \\citep{wedemeyer09}.  Variability of a planet-hosting star could be influenced by the planet itself. Planets orbiting close enough to a star will interact tidally and, possibly, magnetically with the host \\citep{cuntz00}. Magnetic interactions could lead to flares from the reconnection of planetary and stellar fields \\citep{rubenstein00,lanza08}, increased stochastic fluctuations from overall magnetic activity enhancements \\citep{cuntz00}, or periodicities from enhanced plages and faculae surrounding the sub-planetary point on the star \\citep{lanza08,cohen09,kopp11}. Tidal interactions could produce flows and turbulence associated with the tidal bulge \\citep{cuntz00}. They could also spin up the star \\citep{aigrain08}, indirectly increasing overall stellar magnetic activity. These interactions are supported by some evidence (beginning with \\citealt{shkolnik03}), but more definitive conclusions require future, dedicated observations \\citep{lanza11}.  \\subsection{A Selection of Relevant Flare and Variability Studies} \\label{sec:context} There is a long history of research into the frequency and intensity of flares on the Sun and other stars. Especially relevant is recent work by \\citet{hilton10} and \\citet{davenport12}, and references therein, examining large (several $10^4$) samples of M dwarf stars using multi-epoch data in the optical from the Sloan Digital Sky Survey (SDSS) and in the infrared from the Two Micron All Sky Survey (2MASS). \\citet{tofflemire12} specifically assessed the impact of M dwarf flares on exoplanet observations in the infrared using three such stars. Recently, \\citet{kowalski13} conducted a detailed spectrophotometric study in the near UV and optical of 20 M dwarf flares in order to probe the various mechanisms responsible for flare emission. Previous studies in the far and extreme UV are scarcer. \\citet{welsh07} leveraged data from the {\\it Galaxy Evolution Explorer} in the broadband near and far-ultraviolet to find 49 variable sources exhibiting 52 flares. The {\\it Galaxy Evolution Explorer} FUV band data contain the \\siv\\ line we analyzed for variability. In addition, \\citet{mullan06} examined 44 F-M stars in broadband extreme UV time-series data from the {\\it Extreme Ultraviolet Explorer}. The band they utilized is dominated by emission lines of \\ion{Fe}{18} -- \\ion{Fe}{22} formed at coronal temperatures upwards of $10^7$ K, expected in magnetically active regions.  Several previous studies have quantified the stochastic variability of large samples of stars in the optical, most notably employing {\\it Kepler} results to place the Sun's well-characterized variability in the context of other stars (see \\citealt{basri10} and \\citealt{mcquillan12} for examples using {\\it Kepler} data and \\citealt{eyer97} for one using {\\it Hipparcos} data). In addition, it is standard practice to quantify the variability of the exoplanet host star complimentary to radial-velocity or transit measurements, so many individual measurements of stellar variability exist (e.g. \\citealt{dragomir12, kane11, berta11}). However, to the knowledge of the authors this paper presents the first analysis, focusing specifically on the implications for transit observations, of stellar variability in UV line emission flux. ",
        "conclusions": "\\label{sec:discussion}   \\subsection{The Range of Flare Behavior} \\label{sec:flarebehave} \\subsubsection{Strong Flares} Several flares appear in the data that peak at tens of times the quiescent flux. The strongest of these is a flare on Prox Cen, displayed in Figure~\\ref{fig:flarepalette}a  (see also \\citealt{christian04}). This flare raises the continuum-subtracted flux in each band, normalized by the quiescent mean, to values at the peak data point of $27.5\\pm0.2$ in \\cii, $66.7\\pm1.2$ in \\siii, $91.6\\pm0.3$ in \\siv, and $13.5\\pm0.6$ in the continuum. The $\\nfe$ and duration of this flare in each band can be compared with the rest of the flare sample in Figure~\\ref{fig:flaredist}, where these data points are labeled with 39, Prox Cen's dataset number from Table~\\ref{tbl:obs}. Very strong flares also appear in the AD Leo (dataset no. 36; see also \\citealt{hawley03}) and IL Aqr (dataset no. 37; see also \\citealt{france12}) data and are also labeled in Figure~\\ref{fig:flaredist}. The AD Leo flare peaks at $8.2\\pm0.1$ in \\cii, $14.0\\pm0.2$ in \\siii, $33.7\\pm0.1$ in \\siv, and $45.9\\pm0.3$ in the continuum (Figure~\\ref{fig:flarepalette}b). The IL Aqr flare peaks at 15.4$\\pm0.3$ in \\cii, $36.9\\pm0.3$ in \\siii, $46.5\\pm0.3$ in \\siv, and $20.1\\pm1.4$ in the continuum (not included in Figure~\\ref{fig:flarepalette}).  \\subsubsection{Symmetric Flares} Some of the flares in the sample show roughly equal rise and decay times, in contrast to the impulse-decay shape of many flares. The clearest example of such a flare is that which appears in the EG Cha data, plotted in Figure~\\ref{fig:flarepalette}c. Note that this event was not flagged in the continuum or \\cii\\ data because the short span of the data, 0.82 hours, resulted in a poor determination of the quiescent scatter. The continuum and \\cii\\ did not sufficiently exceed the estimated scatter to result in an identification. Other flares that show a relatively clear symmetric photometric profile (though not nearly as clear as that of EG Cha) appear in the \\siv\\ data of EV Lac at mean Julian Date 52172.709 and the AD Leo data at mean Julian Date 51616.119. Many flares in the data appear as though they might be symmetric, but their rise and fall are not adequately resolved.  \\subsubsection{F Star Anomalies} The flare identification algorithm flagged two events on F stars, one in \\cii\\ on $\\delta$ Cep and another in the continuum on $\\beta$ Dor. The latter is displayed in Figure~\\ref{fig:flarepalette}d. Both events are gradual elevations of the flux in a single band relative to the other three. The $\\beta$ Dor flare is particularly curious, as the other three bands clearly do not show the same consistent decline in flux that the elevated continuum flux does. These events seem likely to be true anomalies rather than spurious detections.  \\subsubsection{Multi-peak Flares} Many of the flares in the data exhibit complicated shapes, and might be a superposition of nearby small flares. Figure~\\ref{fig:flarepalette}e depicts a flare on AD Leo that exhibits two peaks, separated by 240 s. This was the clearest specimen of a multi-peak flare. In most such flares, the data are not as clearly resolved above the surrounding scatter and/or different bands show different behavior.  \\subsubsection{Response in Different Lines} We found that flare signals generally vary significantly between bands. In addition, these variations are not consistent from flare to flare. The difference in the response between bands is clear in Figure~\\ref{fig:flarepalette}. For example, the Prox Cen, EG Cha, and AD Leo flares of subplots a, c, and e peak highest in \\siv, with a roughly similar shape in each band. Alternatively, the AD Leo flare depicted in subplot b is the strongest in the continuum, and that of Prox Cen depicted in subplot f is the strongest in \\siii. For the latter, the continuum shows essentially no response, though better signal to noise might reveal otherwise. (Note in Figure~\\ref{fig:flarepalette}f that some \\siii\\ points are negative because the subtracted signal from the Ly$\\alpha$ wing sometimes exceeds the low \\siii\\ signal.) Differences between lines in emitted flux during a flare have also been observed on the Sun (e.g. \\citealt{brekke96}) and other low mass stars (e.g. GJ876, \\citealt{ayres10,france12}).  \\begin{figure} \\includegraphics[width=\\columnwidth]{f14.pdf} \\caption{A palette of flare behaviors, discussed in Section~\\ref{sec:flarebehave}. Each subplot is labeled with the star and the mean Julian Date of $t = 0$. The fluxes are continuum-subtracted, high-pass filtered, and normalized to the mean of the quiescent points. The x symbols denote points flagged for removal before computing excess noise values (\\emph{not} equivalent to the flare duration, see Section~\\ref{sec:flareresults}. Note that the y-axis scales differ.} \\label{fig:flarepalette} \\end{figure}  \\subsection{Risks Flares Pose to Transit Measurements} \\label{sec:flarerisks} All of the detected flares are short lived, with one lasting 27 min, the others lasting less than 20 min, and over half lasting 4 min or less. These could easily be hidden by longer cadence data. Typical cadences in UV exoplanet transit observations are $\\dt \\approx$ 30-60 min, based on recent literature. While hour or longer cadences might often be unavoidable due to instrument or signal-to-noise limitations, a flare hidden in such data could bias a measurement of transit depth (see Section~\\ref{sec:varintro}).  Of the 116 flares we identified (again, excluding FK Com), 57 (roughly one event per 5 h of data) would boost a 60 min integrated flux measurement by $\\gtrsim10\\%$, and all flares (roughly one event per 2.5 h of data) would boost a 60 min flux measurement by $\\gtrsim\\ $1\\%. Flares produced the largest $\\nfe$ values in the \\siv\\ band in roughly 2/3 of the 60~separate events and 7/9 of the events that registered a detection in all four bands. Thus, it appears that transit observations in \\siv\\ are somewhat likely to be more strongly affected by flaring.  The flare stars (namely AD Leo and Prox Cen) account for most of the flares we detected, 92 of 116.  However, 24 flares were identified on objects not classified as ``flare stars.'' Such flares are of particular interest because these stars are more likely to be targeted in exoplanet search programs. These occurred on 5/8 M, 1/8 K, 3/11 G, and 2/5 F stars, and occurred roughly once per 5 h in the time all 32 stars were observed. Of these 24 flares, 9 (roughly one event per 13 h) would boost a 60 min integrated flux measurement by $\\gtrsim10\\%$. Most such flares occurred on M stars (7/9), while the remaining 2/9 are the same event observed in two different bands on the K star EG Cha.  A 10\\% boost in flux exceeds the transit depth in \\cii\\ or \\siii\\ of HD209458b \\citep{madjar04,linsky10}. However, not all stars in the sample are similar in size to HD209458, and it is the relative size of the star and planet that determines the transit depth and thus the impact of a flare. In other words, a strong flare will have less of an impact on a measurement of a Jupiter transiting and M star than a Jupiter transiting an F star. We used the stellar radii from the literature to determine the size of an object that would produce a transit signal of the same amplitude as the boost in flux from a flare in a 60 min integration. These range from 0.02 $R_J$ to 12.5 $R_J$ for the 116 flares identified. Of these, 90 (one flare per 3 h of data) would boost a 60 min flux integration by an amount larger than the signal of an Earth transiting the flaring star, 23 (one flare per 13 h) would boost flux by an amount exceeding a Neptune transit, and 7 (one flare per 42 h) would boost flux by an amount exceeding a Jupiter transit. Limiting the sample to the 24 events on non-flare stars only, 5 flares (one per 31 h) boost flux beyond an Earth signal, 2 (one per 76 h) beyond a Neptune signal, and 1 (one per 153 h) beyond a Jupiter signal.  We did not explore trends in flare rates with respect to stellar properties because the rates were only well constrained on Prox Cen and AD Leo. In addition, the small size, high diversity, and, most importantly, range of flare detection limits in the sample pose problems to such an analysis. Our general conclusion is that all low-mass stars likely pose a risk of flaring near or within a transit observation. Because flares might be so easily hidden in long-cadence data, we recommend that, when possible, transit observers employ minute-scale cadences to inspect their data prior to employing longer cadences for noise suppression. This would enable a sweep of the lightcurve for obvious flare events below the transit timescale. Flares could then be excised and the data, if desired, binned to a longer cadence, correcting for the ``dead time.\" Cadences of arbitrary length are possible with data from the photon-counting detectors common in UV work, provided the data are recorded as a time-tagged event list, rather than time-integrated counts.   \\subsection{Flare Statistics on AD Leo} We detected enough flares on AD Leo to examine their distribution in $\\nfe$. As such, Figure~\\ref{fig:adleoflares} plots the frequency of flares, $\\nu$, with $>\\nfe$ versus $\\nfe$. We fit a power law to these distributions of the form \\begin{equation} \\nu = \\alpha\\nfe^\\beta \\end{equation} using the maximum likelihood method of \\citet{crawford70}. The resulting values of $\\beta$ are $-0.90\\pm0.29$ for \\cii, $-0.92\\pm0.27$ for \\siii, and $-0.82\\pm0.17$ for \\siv. In comparison, previous values include \\begin{itemize} \\item{$-0.82\\pm0.27$ from 21 h of visible and near-ultraviolet observations \\citep{lacy76}} \\item{$-0.62\\pm0.09$ from 111.5 h (spread over $>5$ years) of {\\it U} band observations \\citep{pettersen84}} \\item{$-1.01\\pm0.28$ from $<72$ h of extreme-ultraviolet observations \\citep{audard00}} \\item{$-0.68\\pm0.16$ from 139.7 h of visible observations \\citep{hunt12}.} \\end{itemize} The above values all agree with those we computed for each band. This agreement is consistent with the response of each band tracing common energy deposition events, even for emission resulting from different regions of the stellar atmosphere. To explore further how impulsive energy deposition affects differing regions of a stellar atmosphere, simultaneous, panchromatic flare observations would be desirable.  \\begin{figure} \\includegraphics[width=\\columnwidth]{f15.pdf} \\caption{A cumulative flare-frequency distribution for the flares identified in data for AD Leo. Power-law fits are overplotted. While the data are binned for display, following the methodology of \\citet{crawford70} they were not binned when fitting the power laws.} \\label{fig:adleoflares} \\end{figure}  \\subsection{Size of Detectable Transiting Objects} This work aims to explore the boundaries placed on transit observations by stellar variability in the ultraviolet. Stochastic fluctuations in the host star determine the minimum transit depth that will stand out from these fluctuations, consequently limiting the minimum detectable size of transiting objects. As a metric for this limitation set by each star + band's stochastic fluctuations, we computed the radius of an occulting disk that would produce a transit signal equivalent to $\\s_x$ projected to 3.5~h, $R_{\\s_x}$ -- in essence, the object size needed for a 1-$\\s$ detection of a single transit in the absence of photometric noise. Explicitly, we compute $R_{\\s_x}$ from \\begin{equation} R_{\\s_x}^2 = \\s_x R_\\star^2, \\end{equation} where $R_\\star$ is the radius of the host star. This $R_{\\s_x}$ does not represent an actual detection limit. The true minimum detectable object size depends on the instrument and observing time available. Instead of a true limit, $R_{\\s_x}$ is an instrument independent means of comparing the suitability of stars for transit measurements. This metric is thus free from any assumptions about the number of photons an instrument will collect from the star or what other noise the instrument will add.  Figure~\\ref{fig:disksizes} shows the results, grouped by spectral type. These suggest that, in the absence of photometric uncertainties, roughly Jovian-size disks would produce the smallest detectable transit signal (for a reasonable quantity of data) in a typical system. Indeed, HD209458b, HD189733b, and WASP-12b (see Section~\\ref{sec:intro}) are all of Jovian dimensions. However, $R_{\\s_x}$ spans around an order of magnitude within each spectral type, and more than two orders of magnitude overall.  \\begin{figure} \\plotone{f16.pdf} \\caption{The size in Jupiter radii of an occulting disk that would cause a dip equivalent to the 1-$\\s$ scatter in the stellar flux at $\\dt = 3.5$ h. Filled circles represent confident measurements of $\\s_x$ whereas open circles represent upper limits. Squares represent post-main-sequence stars, circles represent main sequence stars, and diamonds represent WTTS. We group values on the x-axis by spectral type (labels below axis), then by band (symbol color).} \\label{fig:disksizes} \\end{figure}  The results in Figure~\\ref{fig:disksizes} are grouped by spectral type to examine the tradeoff between the smaller $R_\\star$ but higher $\\s_x$ (see Section~\\ref{sec:varcor}) of less massive stars. Smaller stellar disks imply deeper transit signals, but higher levels of stochastic fluctuations better hide these signals. From Figure~\\ref{fig:disksizes}, it appears stellar size trumps $\\s_x$: the smallest stars also permit the smallest objects to produce detectable transits. However, this apparent trend is significantly weakened when the F stars, all post main-sequence giants or sub-giants, are removed.  Of the stars with data robustly sampling their stochastic fluctuations on transit timescales  (mean flux over $5\\sn{-16}$ erg s$^{-1}$ cm$^{-2}$ and total accumulated observations $\\geq3.5$ h), Prox Cen has the most generous $R_{\\s_x}$ limits in each band. These are $1.8\\pm0.1~ R_\\earth$ in \\cii, $<1.5~R_\\earth$ in \\siii, $2.1\\pm0.1~R_\\earth$ in \\siv\\, and $<1.1~R_\\earth$ in the FUV continuum. Interestingly, Prox Cen is classified as a flare star, and the chances that it could flare near or within a transit are significant.  Like Prox Cen, most of the stars in this sample have received flare, variable, or WTTS classifications. The sample is biased: The targets were preferentially selected for variability by the various individuals that commissioned the observations. It seems probable that a less-biased sample would produce detection limits clustered lower than those in Figure~\\ref{fig:disksizes}. Therefore, the prospects for FUV transit measurements of planets the size of Neptune or super-Earths seem promising. Furthermore, if the observed planet hosts an atmosphere inflated to several times the area of the solid disk, like HD209458b \\citep{madjar03}, then in an ideal case even an Earth-size planet might produce a detectable FUV signal within a few folded transits. These prospects are exciting, but it bears remembering that, to attain such limits, photometric noise must be suppressed to below the level of stochastic fluctuations.   \\subsection{Correlations with Stellar Properties} \\label{sec:varcor} We explored correlations between the mean-normalized excess noise measurements in \\cii\\ and \\siv\\ with properties of the sample stars (see Table~\\ref{tbl:props}). Excess noise values in the remaining bands, \\siii\\ and the FUV continuum, were too poorly constrained (essentially, there were too few detections and too many upper limits) to support correlations with any value. As a means of visualizing potential correlations, Figures~\\ref{fig:varcorCII} and~\\ref{fig:varcorSiIV} graph $\\s_x$ values (both detections, black, and upper limits, blue) against each of the stellar properties, excluding mass and luminosity, for the \\cii\\ and \\siv\\ bands respectively.  These figures show that the error bars of the mean-normalized excess noise values and stellar properties often overlap. Thus, changes in the point values even under the 1-$\\s$ error bars could change the Spearman Rank-Order correlation coefficient for the data and, more importantly, confidence that the correlation cannot be explained by randomly scattered points. To account for the uncertainty in how the true values of the points fall, we constructed a Monte-Carlo simulation, generating $10^4$ possible arrangements of the data points given the uncertainties. For each trial in the Monte-Carlo simulation, we randomly drew stellar parameter (e.g., age, $M$, $P_{rot}$) values from Gaussians matching the means and 1-$\\s$ uncertainties. When there was no uncertainty accompanying the measurement we found in the literature, we assigned an uncertainty of 10\\%. For the mean-normalized excess noise values, we did not use Gaussians. Instead we randomly drew values from the actual likelihood distributions of $\\s_x$ for each star/band (Appendix~\\ref{sec:likelihood}).  For each of the $10^4$ such realizations of the data, we recorded probability to exceed (PTE) for the null hypothesis of randomly distributed points (i.e. no correlation) using the Spearman Rank-Order test. We multiplied the PTE by $10^{-4}$ to compute the joint probability of the data producing the correlation coefficient \\emph{and} such a correlation coefficient resulting from random point scatter. The integral of these values over the range of possible correlation coefficients (-1 to 1), provided our overall estimate of the PTE for random scatter in light of the data and uncertainties. We then recorded the significance of the correlation as $1-$PTE.  We carried out this process with and without the $\\s_x$ upper limits and, for both cases, quote the significance of the correlation in Figures~\\ref{fig:varcorCII} and~\\ref{fig:varcorSiIV}. The $\\s_x$ upper limits provide useful constraints on correlations when they fall below the surrounding $\\s_x$ detections. As with the $\\s_x$ detections, when we included upper limits we generated $\\s_x$ values in the Monte-Carlo simulation from the likelihood distribution that produced each upper limit. For instance, an upper limit of 0.01 on $\\s_x$ for a point meant that we randomly drew a value between 0 and 0.01 with roughly uniform probability for each trial in the Monte-Carlo simulation.  The results computed without including the upper limits suggest weak correlations in all cases except $\\s_x$-age and, in \\cii\\, $\\s_x$-$P_{rot}$. The $\\s_x$ upper limits further constrain these correlations, both quantitatively and by eye, and bring \\cii\\ and \\siv\\ into closer agreement. The subsections below address each of these in light of previous literature.  \\begin{figure*} \\plotone{f17.pdf} \\caption{Excess noise versus stellar properties in the \\cii\\ band. Symbols differentiate between spectral types (see legend in bottom right plot). Black symbols represent excess noise detections, while blue symbols represent upper limits. The black numbers above each plot give the significance of the correlation (probability it is not produced by uncorrelated points) using only excess noise detections while the blue number gives the significance including upper limits (see Section~\\ref{sec:varcor}).} \\label{fig:varcorCII} \\end{figure*}  \\begin{figure*} \\plotone{f18.pdf} \\caption{Excess noise versus stellar properties in the \\ion{Si}{4} band, following the same format as Figure~\\ref{fig:varcorCII}.} \\label{fig:varcorSiIV} \\end{figure*}  \\subsubsection{Temperature, Radius, Mass, and Luminosity} \\label{sec:sptcorr} Possible correlations of excess noise with temperature and radius (as well as mass and luminosity, though these are left out of Figures~\\ref{fig:varcorCII} and~\\ref{fig:varcorSiIV}) become strong when upper limits are included. This indicates a more general correlation between excess noise and spectral type: Stars of later spectral type typically exhibit higher $\\s_x$. Correlations between stellar variability and spectral type were explored by \\citet{mcquillan12} for all {\\it Kepler} stars save those with known eclipsing companions (stellar or planetary) and lightcurve discontinuities. They found that cooler, later type stars exhibited high levels of stochastic variability by their metric, in agreement with our findings. Similarly, \\citet{ciardi11} found a relationship between variability and temperature in their analysis of {\\it Kepler} field stars, in both dwarfs and giants. We note, however, that {\\it Kepler} utilizes a broad optical bandpass ($\\sim$4000-9000 \\AA). Therefore, variability in the {\\it Kepler} data reflects processes occurring in the photosphere, rather than the transition region, and there is no guarantee that the two are directly related.  \\subsubsection{Rotation Period, Equatorial Velocity, and Age} Excess noise does not correlate with age, and the hint of a correlation between $\\s_x$ and rotation period is all but eliminated with the inclusion of $\\s_x$ upper limits. A possible anticorrelation seems to exist between $\\s_x$ and \\vsini; however, when $\\s_x$ upper limits are included, the anticorrelation is probably weak enough to dismiss outright. Furthermore, because $v = 2\\pi R/P_{rot}$, the combination of a lack of a $\\s_x$-$P_{rot}$ correlation and the presence of a strong $\\s_x$-$R$ anticorrelation are capable of producing the $\\s_x$-\\vsini\\ trend. Thus, it is the $\\s_x$-$R$ anticorrelation (or, rather, the relationship of $\\s_x$ to spectral type) that drives the $\\s_x$-\\vsini\\ trend. Indeed, $\\s_x$ and the quotient $R/P_{rot}$ (excluding the 8 stars for which we computed $P_{rot}/\\sin{i}$ from \\vsini\\ and $R$) are anticorrelated to roughly the same confidences as $\\s_x$-\\vsini.  Interpreting the lack of correlations with age and rotation in light of previous work is difficult, as we did not find any studies exploring correlations between these properties and stellar stochastic fluctuations. Both age and rotation, however, have been tied to stellar activity as quantified by chromospheric emission (e.g., \\citealt{wilson63,skumanich72,noyes84,mamajek08}). In turn, chromospheric emission likely has a direct relationship with stochastic fluctuations at visible wavelengths \\citep{hall09}. This suggests that younger, faster rotating stars might exhibit higher levels of UV emission line variability.  Given these past results, the lack of $\\s_x$-age and $\\s_x$-$P_{rot}$ correlations in this sample could be explained if either the age-rotation-activity or the variability-activity relationships do not hold for this sample. That either might not hold would not be particularly alarming given that this sample includes stars approaching, dwelling on, and departing the main sequence, whereas the above-mentioned studies only analyze stars on the main sequence. The pre and post main-sequence stars in this sample might exhibit magnetic behavior not in line with that of main-sequence stars, such that relationships between magnetic activity (chromospheric emission) and rotation, age, and/or variability for main-sequence stars are not extensible to this broader sample.  \\subsubsection{Surface Gravity} A strong correlation exists between excess noise and surface gravity in the sample, particularly in \\cii. This is in conflict with the recent results of \\citet{bastien13}, who find an inverse relationship between variability and surface gravity in {\\it Kepler} field stars. However, compared to this work, the stellar sample employed by \\citet{bastien13} is far more restrictive. They limit their study to stars with $4500$ K $< T_{eff}<6500$ K, $2.5< \\log_{10}{g} <4.5$ (cgs units), and relative brightness variations (measured using a 30 min cadence) of $< 0.003$. Applying the same cuts in $T_{eff}$ and $\\log_{10}{g}$ to this stellar sample eliminates any correlation (null hypothesis can produce the observed or stronger correlation with $\\sim$50\\% probability).  The strong correlation present in this data when all stars are included is driven primarily by the F giants, as these have both very low $g$ and $\\s_x$. However, high $\\s_x$, high $g$ M stars also promote the correlation.  \\subsection{Star-Planet Interactions} Close-in planets might interact with their host star magnetically and tidally, producing enhanced activity in the stellar upper atmosphere \\citep{cuntz00}. Evidence for these interactions could be found by monitoring a single star with a transiting planet to search for signs of elevated activity near transit, when the portion of the star interacting with the planet is in view. Indeed, this has been attempted for two systems with transiting planets in this sample, HD209458 \\citep{shkolnik04} and HD189733 \\citep{shkolnik08,fares10}. However, these studies found no clear evidence of enhanced activity correlated with orbital phase for either of these systems.  Alternatively, evidence for star-planet interactions might result from comparing the overall activity or variability of the stars with close-in planets to a sample of control stars. \\citet{shkolnik13} recommended investigations into star-planet interactions in the form of time-resolved observations of UV flux variability rather than activity (i.e. mean line flux), such as those presented in this paper. The small size and high diversity of the 32 stars in this sample that do not host planets make for a poor control. Nevertheless, we looked for increased $\\s_x$ in the six known planet hosts. These hosts are identified in the $\\s_x$ detections and upper limits presented in Figures~\\ref{fig:longvar} and~\\ref{fig:longvar} as open points.  The hosts with $\\s_x$ detections in Figure~\\ref{fig:longvar} do appear clustered at higher $\\s_x$. However, these are all M dwarfs, shown in the previous section to have the highest $\\s_x$ values in the sample. The M dwarf planet hosts do not stand out when grouped with the other M dwarfs in the sample. The G and K planet hosts (HD209458 and HD189733) have insufficient data for a $\\s_x$ detection, allowing only 95\\% upper limits. Because only upper limits are possible, these stars are not constrained to be more variable than other sample stars of the same spectral type. Thus, star-planet interactions, manifested as increased variability, are not supported by the limited volume of this data.    We have examined stellar luminosity fluctuations and stellar flares in a sample of 38 cool stars using 60~s cadence lightcurves constructed from narrow spectroscopic bands containing the \\ion{Si}{3} $\\lm$1206, \\ion{C}{2} $\\lm\\lm$1334,1335, and \\ion{Si}{4} $\\lm\\lm$1393,1402 resonance lines and a combined 36.5 \\AA\\ of interspersed FUV continuum bands. In the high pass filtered lightcurves, we detected 116 flares, occurring roughly once per 2.5 h. Flares commonly radiated more energy in \\siv, relative to quiescent levels, than the other bands. Shorter flares are more prevalent, with over half lasting 4 min or less. Most (90 of 116) of the detected flares could annihilate the signal of an Earth transit, while 7 of the 116 could annihilate the signal of a Jupiter transit. These results highlight the usefulness of minute-scale cadences for finding and removing flares prior to estimating transit depths.  To quantify stochastic fluctuations, we found the maximum-likelihood ``excess noise\" that, in addition to the photometric noise, accounts for the scatter of the high-pass filtered, mean-normalized lightcurves excluding flares. Values of the excess noise, relative to the mean flux, range from about 1-41\\% in \\cii, 8-18\\% in \\siii, 0.9-26\\% in \\siv, and 1-22\\% in the FUV continuum. Where the likelihood distribution of the excess noise was one-sided, we instead quote a 95\\% upper limit on the excess noise. These upper limits on excess noise are often strong enough to be lower than the values for stars where excess noise was detected.  We found significant anticorrelations of excess noise with mass, radius, temperature, and luminosity in \\cii\\ and \\siv. An additional, weaker anticorrelation of excess noise with \\vsini\\ could be an artifact of the strong underlying correlation between radius and rotation rate in the sample. There was no correlation with age or rotation period.  The median levels of stochastic stellar fluctuations we estimated, integrated over a typical transit timescale of 3.5~h, would impose a rough 1-$\\s$ transit detection limit of $\\sim1~R_J$ occulting disks for many of the stars in the sample. However, the range in these limits is broad, spanning from tenths of $R_J$ to tens of $R_J$.  M dwarfs might permit the FUV observation of transiting objects as small as Neptunes or even super-Earths in the absence of photometric noise. While the large fluctuations of some stars might stymie transit spectroscopy in the FUV for any but the largest planets, these results suggest that many planetary host stars might be found for which the limits of FUV transit spectroscopy can be pushed well below the hot Jupiters observed thus far.  {\\it Acknowledgments:} The authors wish to acknowledge the anonymous referee, whose familiarity with the field and resulting thoughtful and constructive comments greatly improved this work. The authors also wish to thank Tom Ayres for his careful reading and helpful suggestions that similarly improved this work. Kevin France acknowledges support through a NASA Nancy Grace Roman Fellowship during a portion of this work. This research has made use of the Exoplanet Orbit Database and the Exoplanet Data Explorer at exoplanets.org \\citep{wright11} as well as extensive use of the SIMBAD database and VizieR catalog access tool, operated at CDS, Strasbourg, France (\\citet{wenger00}, \\citet{ochsenbein00}). The authors wish in particular to acknowledge \\citet{glebocki05} for the excellent stellar rotation catalog that we used in this work. Some of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. This work was supported by NASA grant NNX08AC146 to the University of Colorado at Boulder.  {\\it Facility:} \\facility{{\\it HST} (COS, STIS)}  \\appendix"
    },
    "1402/1402.0245_arXiv.txt": {
        "abstract": "It is suggested that many $\\gamma$-ray bursts (GRBs) are cloaked by an ultra-relativistic baryonic shell that has high optical depth when the photons are manufactured. Such a shell would  not fully block photons reflected or emitted from its inner surface, because the radial velocity of the photons can be less than that of the shell. This avoids the standard problem associated with GRBs that the thermal component should be produced where the flow is still obscured by high optical depth. The radiation that escapes high optical depth obeys the Amati relation. Observational implications may include a) anomalously high ratios of afterglow to prompt emission, such as may have been the case in the recently discovered PTF 11agg, and b) ultrahigh-energy neutrino pulses that are non-coincident with detectable GRB. It is suggested that GRB 090510, a short, very hard GRB with very little afterglow, was an {\\it exposed} GRB, in contrast to those cloaked by baryonic shells. ",
        "introduction": "A longstanding puzzle concerning GRBs % is that their highly super-Eddington luminosities would seemingly imply that the very photons that comprise the GRB should drag out enough baryonic material to obscure themselves. The problem is compounded by the discovery that long GRB occur inside massive stars, meaning that the GRB fireball must plow up material from  its host star as it pushes its way out, not all of which can move out of its way. Neutrons drifting into the fireball ( Eichler \\& Levinson, 1999, Levinson \\& Eichler, 2003), and highly opaque pre-collapse winds from the host stars (Ofek et al., 2013) may also cloak GRBs with an optically thick layer out to $10^{3.5}$ lightseconds.  Naively, this would seem to be enough to completely hide the GRB, which typically lasts less than $10^2$ s.  One suggestion to solve the first problem is that the energy powering the GRB originates on magnetic field lines that thread a black hole that forms at the center of the  collapsing host star (e.g. Levinson and Eichler, 1993), which should be nearly devoid of baryonic material, but this does not address the issue of the plowed-up host star material, which cloaks the fireball even on  horizon-threading field lines.  One solution to the cloaking problem is that many GRBs are seen by observers just off the (angular) edge of the jetted fireball  (Eichler \\& Levinson, 2004, Levinson \\& Eichler 2006), and that this accounts for the Amati and Ghirlanda relations  (Amati, 2002, Ghirlanda  et al 2004 ). The hypothesis successfully predicted an intermediate flat phase for GRB afterglows (Eichler, 2005, Eichler \\& Jontof-Hutter, 2005) that was later confirmed by Swift observations. The \"off-edge observer\" hypothesis also solves the problem\\footnote{See Vurm, Lyubarsky \\& Piran (2013) for a detailed, carefully quantified discussion of this point. They do not, however, consider scattering within the host star envelope off collimating walls, which could lower the photon energies in the observer frame via Compton recoil and pair production.} that because blue-shifted $\\gamma$-ray photospheres should have spectral peaks that are or order $3 \\Gamma kT_{ann}\\gtrsim 1$ MeV, where $kT_{ann}$, the temperature at which pairs annihilate and permit optical transparency, is of order 10 KeV, and $\\Gamma$, the bulk Lorentz factor of the flow, is $\\gtrsim 10^2$,  the spectral peak should then be $\\sim 3$ MeV, whereas most are observed to have spectral peaks at $\\lesssim 300$ KeV.  Taking into account the possibility that pair production opacity limits the photon energy to $\\sim 1$ MeV, there is still some discrepancy with the fact that the distribution of $E_{peak}$ for {\\it long}  GRB (in contrast to short, hard GRB) peaks  comfortably below the pair production threshold. Because the  hypothesis  predicts that the observed spectral peak is softened relative to that seen by a  head-on observer ($\\theta=0$) by a factor of $S(\\theta)\\equiv E_{peak}(0)/E_{peak}(\\theta) ={(1-\\beta \\cos \\theta)\\over(1-\\beta)}=(1 + {\\beta(1-\\cos \\theta)\\over(1-\\beta)})\\sim  1+(\\theta\\Gamma)^2$, where $\\theta$ is the viewing offset angle and $\\Gamma$ is the bulk Lorentz factor, the problem disappears if $ \\theta \\Gamma \\sim 3$. The \"off-edge\" viewing hypothesis, however,  does not yet  explain why {\\it most} GRBs are viewed from such an offset angle, especially considering that the opening  angle for GRB, $\\sim 0.1$, is probably larger than $1/\\Gamma$. (Indeed, about half of all classical GRB, although they show spectral peaks of $\\lesssim 300$ KeV, do  {\\it not } display a  noticeable flat phase in their afterglow.) Neither would it solve the cloaking problem if the cloak has a larger solid angle than the jet, as in the case of a thick stellar wind. ",
        "conclusions": ""
    },
    "1402/1402.5656_arXiv.txt": {
        "abstract": "We study  the oscillations of relativistic stars, incorporating key physics associated with internal composition, thermal gradients and crust elasticity.  Our aim is to develop a formalism  which is able to account for the state-of-the-art understanding of the complex physics associated with these systems.  As a first step, we build models using a modern equation of state including composition gradients and density discontinuities associated with internal phase-transitions (like the crust-core transition and the  point where muons first appear in the core).  In order to understand the nature of the oscillation spectrum, we carry out cooling simulations to provide realistic snapshots of the temperature distribution in the interior as the star evolves through adolescence. The associated thermal pressure is incorporated in the perturbation analysis, and we discuss the presence of $g$-modes arising as a result of thermal effects. We also consider interface modes due to phase-transitions and the  gradual formation of the star's crust and the emergence of a set of shear modes. ",
        "introduction": "Neutron stars are seismically complex, with particular classes of oscillation modes associated with specific parts of the involved physics. This makes the construction of a truly realistic model of the star's oscillation spectrum a complicated task. Nevertheless, we are reaching the point where much of the physics involved is understood at the level required to build moderately realistic models.  The aim of this work is thus quite natural; we want to take important steps towards realism by accounting for processes that are relevant as a neutron star matures. We achieve this by tracking the cooling of a given star from a minute or so after birth through the first several hundred years, paying particular attention to the changes in the thermal pressure and the formation of that star's elastic crust.  The output from our state-of-the-art cooling code provides input for the seismology analysis. The results provide a sequence of snapshots of how the star's oscillation spectrum evolves as the star ages. This is an important advance on previous work in this area.  The problem of relativistic star seismology has been considered since pioneering work by  Thorne \\& Campolattaro in the late 1960s \\cite{thorne_campolattaro_1967}. Like much of the subsequent work, the initial focus was on the mathematical formulation of the problem with the complex supranuclear physics playing a secondary role. Thus, the bulk of the literature is focussed on perfect fluid stars, often without any consideration of the interior composition and state of matter.  In fact, many studies have made use of rather ad hoc polytropic models which only capture the rough properties of a realistic equation of state. Nevertheless, this has led to an understanding of the basic nature of the stellar spectrum, like the fundamental $f$-mode and the pressure restored $p$-modes \\cite{lindblom_detweiler_1983, detweiler_lindblom_1985}, and some discoveries, like the existence of the gravitational-wave $w$-modes, \\cite{kokkotas_schutz_1992}. However, the understanding of the fine details of the problem has so far been developed one piece at a time. Gravity $g$-modes arising because of composition gradients have been considered \\cite{finn_1986, finn_1987, strohmayer_1993, miniutti_etal_2003}, the crust elasticity has been accounted for (especially for shear modes) \\cite{schumaker_thorne_1983}, the roles of core superfluidity \\cite{comer_etal_1999, andersson_etal_2002, lin_etal_2008} and the star's magnetic field \\cite{sotani_etal_2008, colaiuda_etal_2009, gabler_etal_2012, gabler_etal_2013} have been investigated and oscillations of proto-neutron stars and their detectability~\\cite{ferrari_etal_2003} have been studied, as well. In the last decade, much attention has been focussed on the Coriolis restored $r$-modes~\\cite{andersson_1998, andersson_kokkotas_2001, lockitch_etal_2001, haskell_etal_2009}, which are relevant because they may be driven unstable due to the gravitational radiation they generate. In this context, much of the complex microphysics related to dissipative processes has been discussed, progressing our understanding of the involved processes considerably.  Neutron star oscillations may impact on a range of observations, involving in particular radio and X-ray timing and gravitational waves. At the present time, the most promising connection between observations and theory is provided by the observed quasiperiodic oscillations seen in the X-ray tail of magnetar giant flares \\cite{watts_strohmayer_2006, samuelsson_andersson_2007}. The inferred oscillation frequencies match those expected for the elastic crust reasonably well, and provide the first credible example of actual neutron star asteroseismology.  There are, of course, issues to be resolved, especially concerning the dynamical role of the interior magnetic field in these systems \\cite{colaiuda_kokkotas_2011}. A more indirect example comes from X-ray timing of fast spinning accreting neutron stars in Low-Mass X-ray Binaries \\cite{strohmayer_mahmoodifar_2014, andersson_etal_2014}.  The spin of these stars appears to be limited by some process. One of the leading contenders for the underlying mechanism is the $r$-mode instability which leads to the star losing angular momentum at a rate that could balance the spin-up due to accretion \\cite{andersson_etal_1999, bondarescu_etal_2007}. Finally, as we are getting closer to the advent of gravitational-wave astronomy, there has been searches for neutron star oscillations from a range of systems \\cite{ligo_search1, ligo_search2, ligo_search3, ligo_search4}. At present, these studies may only be providing relatively uninteresting upper limits on the possible signals. With an advanced generation of detectors coming online in the next few years, this could turn to actual detections.  The existing body of work allows us to piece together a picture of neutron star dynamics. This understanding should be valid provided the star is in a regime where the different ingredients remain distinct. Unfortunately, this is unlikely to be the case. Hence, there is a pressing need to develop a new generation of models that account for as much of the relevant physics as possible. The present work should be seen in that context.  The physics associated with realistic neutron star dynamics is daunting. Many aspects are  rather poorly understood, in particular concerning the deep core (at several times the nuclear saturation density). Nevertheless, a focussed effort on this problem is timely. First of all, gravitational-wave astronomy should become reality in the next few years. This promises to provide us with observational data (perhaps most likely from the merger of compact binaries) that need to be matched against the best possible theoretical models. Secondly, our understanding of the principles associated with neutron star superfluidity has improved considerably in the last decade. In particular, we have constraints on the superfluid parameters (the superfluid pairing gaps or, equivalently, the critical temperatures) from the observed real-time cooling of the remnant in Cassiopeia A \\cite{shternin_etal_2011}. In order to be able to make productive use of future observational results, we need to improve our level of modelling.  The basic requirement for progress on this problem is a realistic equation of state which accounts for the two-fluid nature of a neutron star's core, including information about the superfluid gap energies and the entrainment parameters. We also need a comprehensive formalism for modelling dynamics which accounts for all the desired parameters. Progress on the first of these issues was recently made by Chamel \\cite{chamel_2008}, who provided the first ever consistent equation of state including entrainment. The second problem has been explored in basic models \\cite{andersson_etal_2013}, to the point where there are no technical stumbling blocks preventing more realistic studies.  In this work we  take the first steps towards true realism, starting from the classic fluid formalism of Detweiler \\& Lindblom \\cite{detweiler_lindblom_1985} and extending it to account for density discontinuities associated with distinct phase-transitions, interior composition gradients, thermal pressure, and the elastic crust which will form and grow in thickness as the neutron star cools. We use this model  to study the evolution of the star's oscillation modes as it matures. To do this we couple the oscillation mode calculation to the long-term cooling of the star. At different points in the thermal evolution we output the temperature profile and feed it into the mode calculation.  The results of this exercise shed light on the influence of thermal effects on the various oscillation modes of the star. As expected, we find that the fundamental $f$-mode and the various $p$-modes are only weakly affected by the changes in temperature. Meanwhile, the gravity $g$-mode spectrum changes completely as the thermal pressure weakens and we demonstrate how only a few interface modes (arising from density ``discontinuities\") remain (in the considered frequency range above 10~Hz or so)  when the star reaches maturity. Finally, the presence of the elastic crust enriches the spectrum by shear modes, which also evolve as the star cools and the crust region grows.  This paper is the first in a series where we aim to build a truly realistic model of a dynamical neutron star. Our initial focus is on issues relating to composition and thermal effects. This will lead naturally on to the role of superfluidity, which will be considered in a follow-up paper. The basic reason for the division is that the problems we consider here can be modelled within the standard single-fluid framework, while superfluidity adds  degrees of freedom that make the problem richer (and obviously more complicated). The same is true for the star's magnetic field, which would eventually have to be included in the model (involving issues associated with the presence of a superconducting core \\cite{glampedakis_etal_2011}).  The layout of the paper is as follows: Section \\ref{sec:ns_model} is devoted to the physics which we take into account for the background model; in Section \\ref{sec:perturbations}, we discuss the perturbation equations and explain how these are modified to account for additional physics (the perturbation equations for the elastic crust are given in Appendix \\ref{sec:app_eq_crust_polar}). A general numerical strategy for solving the perturbation equations in the interior of a multi-layered star is presented in Section \\ref{sec:num_strategy} and Appendix \\ref{sec:app_numerics}; Section \\ref{sec:num_strategy} also covers the reasoning for a new set of equations for perturbations of the perfect fluid, which aids in solving the perturbation problem at low frequencies; the actual equations can be found in Appendix~\\ref{sec:app_eq_low_freq}. Finally, Section \\ref{sec:results} contains the results and Section \\ref{sec:summary} summarises our work.  Unless stated otherwise, we use units in which $G=c=1$ and Misner, Thorne and Wheeler (MTW) \\cite{MTW1973} conventions throughout the paper. ",
        "conclusions": "\\label{sec:summary}  The aim of this paper was to establish a formalism that allows us to study neutron star oscillations in full general relativity, accounting for as much realistic microphysics as possible. So far we extended the standard perfect fluid formalism to account for thermal pressure in the star's core and  derived a new set of equations which govern perturbations in the solid crust. We provided sample results for a particular realistic equation of state. However, our numerical code is generic in the sense that we can easily use a different equation of state straightaway.  We focus our attention on the low-frequency part of the oscillation spectrum (above 15~Hz or so due to computational limitations). In this regime we find a set of interface modes, firstly due to the (artificial) density discontinuity at the crust-core interface and secondly there are interface modes associated with the crust region, due to sharp changes in the low-density equation of state.  The composition gradient shifts the frequency of the artificial interface mode to slightly higher frequencies, and we find a set of composition $g$-modes in the low frequency regime. The frequencies of these $g$-modes are slightly lower than literature values from pervious studies---this is because the composition gradient is rather small throughout nearly the entire core; only close to the crust-core transition is the composition gradient reasonably pronounced.  When we account for thermal pressure due to neutrons and protons in the core, we find that a number of thermal $g$-modes enter the low frequency part of the spectrum. Meanwhile, the high frequency interface modes are shifted to even higher frequencies whereas the composition $g$-modes stay unaffected. Tracking the thermal evolution of the neutron star for the first few hundred years, we investigate how the frequencies of the various modes evolve as the star matures through adolescence. After approximately 10 years, all thermal $g$-modes have dropped below $15\\unit{Hz}$ and the thermal effect on the interface modes has almost vanished entirely. During the next 100 years, as the star continues to cool, the frequencies of the interface modes change slightly, finally leaving us with the spectrum of a cold star.  Finally, we considered the crystallization of the crust and investigated the associated changes in the spectrum. From our thermal evolution, we determined  the solid region in the outer parts of the star.  We quantified how the crust elasticity affects the interface modes associated with the outer layers of the star and we discovered that they are suppressed by shear stress. We discussed how the spectrum is enriched by shear modes. As expected, our results for the shear modes are in good agreement with results from previous work \\cite{mcdermott_etal_1988}.  In summary, we have taken the first step towards a comprehensive computational technology to study quasinormal modes of realistic compact stars. We plan to extend the model to account for a multifluid core where the neutrons are superfluid and the protons are superconducting. In doing so we will consider recent equation of state data, in particular concerning the entrainment effect, both in the core and the crust. We will also update the cooling sequence to account for superfluid effects. We expect to report on results in these directions in the not too distant future."
    },
    "1402/1402.2289_arXiv.txt": {
        "abstract": "Direct $N$-body simulations of globular clusters in a realistic Milky Way-like potential are carried out using the code \\texttt{NBODY6} to determine the impact of the host galaxy disk mass and geometry on the survival of star clusters. A relationship between disk mass and star cluster dissolution timescale is derived. These $N$-body models show that doubling the mass of the disk from $5\\times 10^{10}$~M$_{\\odot}$ to $10\\times10^{10}$~M$_{\\odot}$ halves the dissolution time of a satellite star cluster orbiting the host galaxy at 6 kpc from the galactic center. Different geometries in a disk of identical mass can determine either the survival or dissolution of a star cluster orbiting within the inner 6 kpc of the galactic center. Furthermore, disk geometry has measurable effects on the mass loss  of star clusters up to 15 kpc from the galactic center. $N$-body simulations performed with a fine output time step show that at each disk crossing the outer layers of star clusters experience an increase in velocity dispersion of $\\sim$5\\% of the average velocity dispersion in the outer section of star clusters. This leads to an enhancement of mass-loss -- a clearly discernable effect of disk shocking. By running models with different inclinations we determine that star clusters with an orbit perpendicular to the Galactic plane have larger mass loss rates than both clusters evolving in the Galactic plane or in an inclined orbit. ",
        "introduction": "In the current paradigm of galaxy formation smaller structures merge into larger ones from the Big Bang up to the present day (White \\& Rees 1978).  Galaxies grow through two main processes: the hierarchical merging with smaller galaxies and the accretion of fresh gas fuelling new star formation. These different mechanisms contribute to the growth of the disk, bulge and halo.  As the first significant stellar structures to form, globular clusters witness the entire evolution of their host galaxy as satellite systems. Indeed, globular clusters are believed to  follow galaxies during galaxy mergers and close encounters (e.g.\\ West et al.\\ 2004). The gravitational potential of their host galaxy has a direct influence on the survival of globular clusters: they lose stars through tidal stripping and disk shocking. In turn, these lost stars contribute to the build-up of the galaxy's stellar halo. The evolution of host galaxy and globular clusters are clearly connected.  The aim of this work is to derive, using numerical simulations, the importance of the host galaxy disk mass and size on the evolution and survival of star clusters. How are satellite stellar systems affected by disks and bulges of changing mass and with different geometries? And reciprocally, how do satellite stellar systems contribute, through their dissolution, to the formation of the halo of the host galaxy?  D'Onghia et al.\\ (2010) show that halo and disk shocking efficiently deplete the satellite population of dark matter halos within 30 kpc of the Milky Way center. The results of that study cannot be directly applied to globular clusters because of their much higher densities compared to dark matter halos.  In their important paper, Gnedin \\& Ostriker (1997) model the dynamical evolution of the Galactic globular cluster system using a Fokker-Planck code. These authors build on earlier analytical work by Aguilar, Hut \\& Ostriker (1988) and Kundic \\& Ostriker (1995) among others. Gnedin \\& Ostriker (1997) give analytical expressions to estimate the impact of disk shocking and also call for numerical simulations to be carried out.  With the recent progress in computational capacity, large $N$-body simulations can now be carried out in order to determine the physical mechanisms that govern the dynamical evolution of globular clusters in a galactic potential. Recently,  Renaud \\& Gieles (2013)  found that after a merger event of the host galaxy the mass loss rate of a satellite star cluster increases. Interestingly, Renaud \\& Gieles (2013) also find that even if the tidal forces reach a maximum during the merger itself they are too short-lived to have a significant impact on the long term survival of star clusters.  Previous work on the mass loss of star clusters focused on internal dynamical effects and stellar evolution. Vesperini \\& Heggie (1997) carried out numerical simulations of star clusters and determined their mass loss rates through a Hubble time. Baumgardt \\& Makino (2003) established an analytical dissolution time scale for star clusters and showed that one third of the cluster mass is lost due to stellar evolution alone. A recent review by of recent $N$-body studies can be found in Portegies Zwart et al.\\ (2010).  The approach in our work is purely numerical and different in nature to Gnedin \\& Ostriker (1997) given  that no assumptions or explicit expressions to treat the impact of the disk are used. $N$-body  models of star clusters were run where the properties of the star cluster remain identical but the mass and geometry of the galactic disk change. The effect of the disk on the star cluster is computed as a part of the numerical calculations of the gravitational force experienced by each star.  The above is carried out with the code \\texttt{NBODY6} used for the study of the dynamics of star clusters through $N$-body simulations. \\texttt{NBODY6} now includes a detailed model of the host galaxy where star clusters evolve as satellites (Aarseth 2003). The current set-up of \\texttt{NBODY6} includes the tools to model a Milky-Way type galaxy with three distinct components: disk, bulge, and halo. The gravitational force for each star of the cluster is computed at each time step by taking into account the effect of all other stars, and of the disk, bulge, and halo. We make use of this new capacity of \\texttt{NBODY6} to run several models of star clusters where the mass and physical size of the disk are different between models.  Throughout this work, a Hubble time of 13.5 Gyr is adopted, in agreement with the results of the Wilkinson Microwave Anisotropy Probe (Spergel et al.\\ 2003). ",
        "conclusions": ""
    },
    "1402/1402.5526_arXiv.txt": {
        "abstract": "{We present a study of the 3-dimensional environment for a sample of 386 galaxies in the \\textbf{C}atalogue of \\textbf{I}solated \\textbf{G}alaxies (CIG, Karachentseva 1973) using the Ninth Data Release of the Sloan Digital Sky Survey (SDSS-DR9).} {We aim to identify and quantify the effects of the satellite distribution around a sample of galaxies in the CIG, as well as the effects of the Large Scale Structure (LSS).} {To recover the physically bound galaxies we first focus on the satellites which are within the escape speed of each CIG galaxy. We also propose a more conservative method using the stacked Gaussian distribution of the velocity difference of the neighbours. The tidal strengths affecting the primary galaxy are estimated to quantify the effects of the local and LSS environments. We also define the projected number density parameter at the 5$^{\\rm th}$ nearest neighbour to characterise the LSS around the CIG galaxies.} {Out of the 386 CIG galaxies considered in this study, at least 340 (88\\% of the sample) have no physically linked satellite. Following the more conservative Gaussian distribution of physical satellites around the CIG galaxies leads to upper limits. Out of the 386 CIG galaxies, 327 (85\\% of the sample) have no physical companion within a projected distance of 0.3\\,Mpc. The CIG galaxies are distributed following the LSS of the local Universe, although presenting a large heterogeneity in their degree of connection with it. When present around a CIG galaxy, the effect of physically bound galaxies largely dominates (usually by more than 90\\%) the tidal strengths generated by the LSS.} {The CIG samples a variety of environments, from galaxies with physical satellites to galaxies with no neighbours within 3\\,Mpc. A clear segregation appears between early-type CIG galaxies with companions and isolated late-type CIG galaxies. Isolated galaxies are in general bluer, with likely younger stellar populations and rather high star formation with respect to older, redder CIG galaxies with companions. Reciprocally, the satellites are redder and with an older stellar populations around massive early-type CIG galaxies, while they have a younger stellar content around massive late-type CIG galaxies. This suggests that the CIG is composed of a heterogeneous population of galaxies, sampling from old to more recent, dynamical systems of galaxies. CIG galaxies with companions might have a mild tendency (0.3-0.4 dex) to be more massive, and may indicate a higher frequency of having suffered a merger in the past.} ",
        "introduction": "\\label{Sec:intro}  Isolated galaxies are located, by definition, in low-density regions of the Universe, and should not be significantly influenced by their neighbours. Does a separate population of isolated galaxies exist, or are isolated galaxies simply the least clustered galaxies of the Large Scale Structure (LSS)? It is assumed that over the past several billion years the evolution of these objects has largely been driven by internal processes. A significant population of isolated galaxies is of great interest for testing different scenarios of the origin and evolution of galaxies. In this sense, isolated galaxies are an ideal sample of reference for studying the effects of environment on different galaxy properties. Such a sample would represent the most nurture-free galaxy population.  Studies of isolated galaxies can be argued to begin with the publication of the Catalogue of Isolated Galaxies \\citep[CIG;][]{1973AISAO...8....3K}. The AMIGA (\\textbf{A}nalysis of the interstellar \\textbf{M}edium of \\textbf{I}solated \\textbf{GA}laxies\\footnote{\\texttt{http://amiga.iaa.es}}) project \\citep{2005A&A...436..443V} is based upon a re-evaluation of the CIG. It is a first step in trying to identify and better characterise isolated galaxies in the local Universe. \\citet{2005A&A...436..443V} argued that 50\\% or more galaxies in the CIG show a homogeneous redshift distribution. \\citet{2006A&A...449..937S}, and more recently \\citet{2012A&A...540A..47F}, found that 2/3 of the CIG are Sb-Sc late-type galaxies, and 14\\% are early-type. This implies an extremely high late-type fraction and extremely low early-type population. At intermediate redshift, \\citet{2012MNRAS.419.3018C} found that early-type systems in higher density regions tend to be more extended than their counterparts  in low density environments. Taking into account the effect of the local environment, \\citet{2013MNRAS.434..325F} show that the number of satellites around a galaxy affects its size. CIG galaxies have larger sizes than galaxies in the \\citet{2010ApJS..186..427N} sample with zero or one satellite, which are also larger than galaxies in \\citet{2010ApJS..186..427N} with two or more satellites.  The distribution of satellites (faint companions) around isolated primary galaxies provides important information about galaxy formation, as well as a critical test of the $\\Lambda$CDM model on small scales \\citep{1987MNRAS.226..543E,2007AAS...21112602C,2010ApJ...709.1321A,2012MNRAS.425.2817F,2013JCAP...03..014A,2013ApJ...772..109B}. This explains the growing interest for studying the satellite distribution \\citep{2003ApJ...598..260P,2005MNRAS.356.1045S,2012MNRAS.427..428G}, and for exploring the link between galaxy properties and the satellite population \\citep{2007ApJ...658..898P,2011MNRAS.417..370G,2011AstBu..66..389K,2012MNRAS.424.1454E,2013ApJ...770...96G}.  According to a previous study \\citep{2013A&A...560A...9A}, the criteria proposed by \\citet{1973AISAO...8....3K} to remove fore- and background galaxies are not fully efficient. About 50\\% of the neighbours, considered as potential companions, have very high recession velocities with respect to the central CIG galaxy: the condition is too restrictive, and may consider as not isolated galaxies slightly affected by their environment. On the other hand, about 92\\% of neighbour galaxies showing recession velocities similar to the corresponding CIG galaxy are not considered as potential companions by the CIG isolation criteria, and may have a non negligible influence on the evolution of the central CIG galaxy. This motivates us to extend the study, taking into account nearby and similar redshift companions to identify physical satellites affecting the evolution of the central CIG galaxy, so as to provide a more physical estimation of the isolation degree of the CIG. About 60\\% of the CIG galaxies have no major (similar-size) companion in the SDSS, according to the CIG (purely photometric) isolation criteria. Nevertheless, considering the third dimension, only 1/3 of the sample has no similar redshift neighbours \\citep{2013A&A...560A...9A}.  In this context the CIG represents an excellent sample to study the relation of galaxy properties on both local and large-scale environments.  In the present work, we aim to identify and quantify the effects of the satellite distribution around a sample of CIG galaxies, as well as the effects of the Large Scale Structure. This study is organised as follows: in Sect.~\\ref{Sec:data}, the sample and the data used are presented. The method to identify the potential satellite galaxies is described in Sect.~\\ref{Sec:physical}. In Sect.~\\ref{Sec:isolparam}, we describe the parameters used to quantify the environment. We present our results in Sect.~\\ref{Sec:results} and the associated discussion in Sect.~\\ref{Sec:discussion}. Finally a summary and the main conclusions of the study are presented in Sect.~\\ref{Sec:con}. Throughout the study, a cosmology with $\\Omega_{\\Lambda 0} = 0.7$, $\\Omega_{\\rm{m} 0} = 0.3$, and $H_{0}=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ is assumed. ",
        "conclusions": "\\label{Sec:con}  We present a study of the 3-dimensional environment for a sample of 386 galaxies in the \\textbf{C}atalogue of \\textbf{I}solated \\textbf{G}alaxies \\citep[CIG;][]{1973AISAO...8....3K}, using the Ninth Data Release of the Sloan Digital Sky Survey (SDSS-DR9).  We identify and quantify the effects of the satellite distribution around a sample of galaxies in the CIG. To recover the physical satellites around the CIG galaxies, we first focus on the satellites which are within the escape speed of each CIG galaxy. We also propose a more conservative method based on the stacked Gaussian distribution of the velocity difference of the neighbours, which gives an upper limit to the influence of the local environment.  In comparison to a previous study \\citep{2013A&A...560A...9A} we can estimate the effect of the physical associations that were not taken into account by the CIG isolation criteria, which could also have a non negligible influence on the evolution of the central CIG galaxy. The tidal strengths affecting the primary galaxy are estimated to quantify the effects of the local and Large Scale Structure (LSS) environments. To characterise the LSS around the CIG galaxies, we define the projected number density parameter at the 5$^{\\rm th}$ nearest neighbour.  Our main conclusions are the following:  \\begin{enumerate} \\item Out of the 386 CIG galaxies considered in this study, at least 340 (88\\% of the sample) have no physically linked satellite. Following the more conservative Gaussian distribution method to identify physical satellites around the CIG galaxies leads to upper limits: out of the 386 CIG galaxies, 327 galaxies (85\\% of the sample) would have no physical companion within a projected distance of 0.3\\,Mpc.  \\item Consequently, about 12\\% (and up to 15\\%) of the CIG galaxies have physically bound satellites. CIG galaxies with companions might have a mild tendency (0.3-0.4 dex) to be more massive, which suggests a higher frequency of having suffered a merger in the past. Satellites are in general redder, brighter, and bigger for more massive central CIG galaxies. Also, massive elliptical and lenticular CIG galaxies tend to have satellites with earlier types than similar mass spiral CIG galaxies.  \\item Although 15\\% at most of the CIG galaxies in the sample have physically bound satellite, almost all galaxies (97\\%) can be directly related to a LSS. The very large scatter in the quantification of the LSS shows that the CIG includes both galaxies dominated by their immediate environment as well as galaxies almost free from any external influence.  \\item The continuous distributions of the $\\eta_{k, \\rm LSS}$ and $Q_{\\rm LSS}$ isolation parameters show that the CIG spans a variety of environments. The connection of the CIG galaxies with the LSS is obvious due to the excess of similar redshift galaxies between 0.3 and 3\\,Mpc. The CIG galaxies are distributed following the LSS of the local universe, although presenting a large heterogeneity in their degrees of connection with it.  \\item To evaluate the role of the physically bound satellites with respect to the large scale environment, we compare the magnitudes of the sum of the tidal strengths produced by the physical companions to the sum of the tidal strengths created by all the galaxies in the LSS. When at least one physical companion is present near a CIG galaxy, its effect largely dominates (usually more than 90\\%) the tidal strengths generated by the LSS.  \\item To delimit the role of the environment on the physical properties of the galaxies we compare, within the CIG, the most isolated galaxies to the galaxies with companions. We find a clear segregation between CIG galaxies with companions and isolated CIG galaxies. Isolated galaxies are in general bluer, with a younger stellar population and rather high star formation with respect to the older, redder galaxies with companions. These $(u-r)$ colours, in combination with the morphological trends and the $(g-r)$ colours, lead to a coherent view where isolated star forming galaxies are separated from older elliptical galaxies with companions.  \\item Conjointly, we find that the satellites are redder and with older stellar populations around massive early-type CIG galaxies while they have a younger stellar content around massive late-type CIG galaxies. This means that if the local environment has an influence on the evolution of the CIG galaxies, reciprocally, the satellites around the CIG galaxies may also be affected by the nature of the primary galaxy. This suggests that the CIG is composed of a heterogeneous population of galaxies, sampling old systems of galaxies but also spanning more recent, dynamical systems of galaxies.  \\item As mentioned, the CIG samples a variety of environments, from galaxies in interaction with physical satellites to galaxies with no neighbours in the first 3\\,Mpc around them. Hence, in the construction of catalogues of galaxies in relation to their environments (isolated, pairs, triplets, groups of galaxies), redshift surveys are required in order to distinguish small, faint, physically bound satellites from a background projected galaxy population and reach a more comprehensive 3-dimensional picture of the surroundings.  \\end{enumerate}"
    },
    "1402/1402.6713_arXiv.txt": {
        "abstract": "The conservation of energy, linear momentum and angular momentum are important drivers for our physical understanding of the evolution of the Universe. These quantities are also conserved in Newton's laws of motion under gravity \\citep{Newton:1687}. Numerical integration of the associated equations of motion is extremely challenging, in particular due to the steady growth of numerical errors (by round-off and discrete time-stepping, \\cite{1981PAZh....7..752B,1993ApJ...415..715G,1993ApJ...402L..85H,1994LNP...430..131M}) and the exponential divergence \\citep{1964ApJ...140..250M,2009MNRAS.392.1051U} between two nearby solution. As a result, numerical solutions to the general N-body problem are intrinsically questionable \\citep{2003gmbp.book.....H,1994JAM....61..226L}.  Using brute force integrations to arbitrary numerical precision we demonstrate empirically that ensembles of different realizations of resonant 3-body interactions produce statistically indistinguishable results. Although individual solutions using common integration methods are notoriously unreliable, we conjecture that an ensemble of approximate 3-body solutions accurately represents an ensemble of true solutions, so long as the energy during integration is conserved to better than 1/10.  We therefore provide an independent confirmation that previous work on self-gravitating systems can actually be trusted, irrespective of the intrinsic chaotic nature of the N-body problem. ",
        "introduction": "Newton's law of gravitation is one of the fundamental laws in the Universe that holds everything together. Although formulated in the 17th century, scientists today still study the consequences, in particular those of many-body systems, like the solar system, star clusters and the Milky Way galaxy.  General analytic solutions to the N-body problem only exist for configurations with one mass, commonly referred to as $N=1$ solutions, and for two masses (equivalently named $N=2$, \\cite{Kepler:1609, Newton:1687}). Problems for $N\\rightarrow \\infty$ can be reduced via Liouville's theorem for Hamiltonian systems to the collisionless Boltzmann equation \\citep[][but see also \\cite{1860dMaxwell}]{1868Bolzmann,1968SvPhU..10..721V}, and therefore analytic solutions for the global distribution function exist.  Solutions for $N$ in between these two limits are generally realized by computer simulations. These so-called $N$-body simulations have a major shortcoming in that the solution to any initial realization can only be approximated.  The main limiting factors in numerically obtaining a true solution include errors due to round-off and approximations both in the integration and in the time-step strategy \\citep{1996magr.meet..167K}.  These generally small errors are magnified by the exponentially sensitive dependence on the 6N-dimensional phase-space coordinates, position and velocity \\citep{1964ApJ...140..250M}. As a consequence, the solution for a numerically integrated self-gravitating system of $N$ masses diverges from the true solution. This error can be controlled to some degree by selecting a phase-space volume preserving or a symplectic algorithm \\citep{1991AJ....102.1528W} and by reducing the integration time step \\citep{1975ARA&A..13....1A,2008MNRAS.386..295H}. The latter however, cannot be reduced indefinitely due to the accumulation of numerical round-off in the mantissa, which is generally limited to 53 bits (64 bits in total, but 11 bits are reserved for the exponent, resulting in only about 15 significant digits).  The exponential divergence subsequently causes this small error to propagate to the entire system on a dynamical time scale \\citep{1993ApJ...415..715G}, which is the time scale for a particle to cross the system once. The result of these errors together with the exponential divergence, is the loss of predicting power for a numerical solution to a self gravitating system with $N>2$ after a dynamical time scale. One can subsequently question the predicting qualities of $N$-body simulations for self-gravitating systems, and therewith their usefulness as a scientific instrument.  We address this question for $N=3$ by brute-force numerical integration to arbitrary precision.  The choice of $N=3$ is motivated by the realization that this represents the first fundamental irregular configuration with the smallest possible number of objects that cannot be solved analytically and cannot be addressed with collisionless theory.  In addition, 3-body encounters form a fundamental and frequently occurring topology in any large $N$-body simulation, and therefore also drive the global dynamics of these larger systems. ",
        "conclusions": "The properties of the binary and the escaper of a three-body system can be described in a statistical way. This is consistent with the findings in previous analytic \\citep{1976MNRAS.176...63M} and numerical \\citep{2004ASPC..316...45V} studies.  This behavior was named quasi-ergodicity by \\cite{1976MNRAS.176...63M}. We confirm that this behavior remains valid also for converged 3-body solutions.  Based on the symmetry of the distribution in dissolution times (see Fig.\\,\\ref{Fig:CumulativeError}), the final parameters of the binary and escaper, as well as the consistency of the mean and median values of the inaccurate simulations when compared to the converged solution (see Fig.\\,\\ref{Fig:CumulativeError} and Fig.\\,\\ref{Fig:fdissolved}) we argue that global statistical distributions are preserved irrespective of the precision of the calculation as long as energy is preserved to better than 1/10th of the initial energy of the system. Although we have tested only three algorithms for solving the equations of motion we conjecture that the statistical consistency may be preserved also for some other direct $N^2$ methods, and these may also require that energy and angular momentum are preserved to $\\leq 1/10$th.  If such direct $N$-body methods comply to the same statistical behavior for collision-less ($N\\gg 3$) systems, it will be interesting to investigate how also other --non-$N^2$-- algorithms, like the hierarchical tree-method \\citep{1986Natur.324..446B} or particle-mesh methods \\citep{Hockney1988} behave in this respect.  In studies of self-gravitating systems which adopt the 4th order Hermite integrator, energy and angular momentum are generally conserved up to $\\aplt 10^{-6}$ per dynamical time. Only those simulations in which this requirement is met are often considered reliable and suitable for scientific interpretation.  Proof for this seemingly conservative choice has never been provided, and it is unknown whether or not the numerical error and the exponential divergence are not preventing certain parts in parameter space to be accessed, or new physically inaccessible parts in parameters space being explored.  We argued that for the resonant 3-body problem the error made during the integration of the equations of motion poses no problem for obtaining scientifically meaningful results so long as energy is conserved to better than about one-tenth of the initial total energy of the system.  In that case resonant 3-body interactions should be treated as an ensemble average, and individual results only contribute statistically.  By means of numerical integration until a converged solution is obtained we find that the statistical properties of the binary and the escaper resulting from a 3-body resonant encounter are deterministic. This behavior is not guaranteed to propagate to larger $N$ (see also \\cite{1992MNRAS.259..505Q}); $N>3$ requires independent testing, because these introduce more complex solutions in the form of, for example binary-binary outcomes and hierarchical triples.  The more extended parameter space for increasing $N$ from 3 to $N=4$ is quite dramatic, in particular for solving the system until a converged solution is reached.  {\\bf Acknowledgements} We would like to thank Douglas Heggie for many in-depth discussions, but also Alice Quillen, Piet Hut, Jun Makino, Steve McMillan, Vincent Icke and Inti Pelupessy for discussions and comments on the manuscript, as well as the anonymous referee for careful reading and detailed comments. This work was supported by the Netherlands Research Council NWO (grants \\#643.200.503, \\#639.073.803 and \\#614.061.608) and by the Netherlands Research School for Astronomy (NOVA).  Part of the numerical computations were carried out on the Little Green Machine at Leiden University and on the LISA cluster at SURFSara in Amsterdam."
    },
    "1402/1402.1075_arXiv.txt": {
        "abstract": "A significant portion of transients measured by spacecraft at 1~AU does not show the well-defined properties of magnetic clouds (MCs). % Here, we propose a new class of complex, non-MC ejecta resulting from the interaction of two CMEs with different orientation, which differ from the previously studied multiple-MC event. At 1~AU, they are associated with a smooth rotation of the magnetic field vector over an extended duration and do not show clear signs of interaction. We determine the characteristics of such events based on a numerical simulation and identify and analyze a potential case in the long-duration CME measured {\\it in situ} in 2001 March 19--22. Such events may result in intense, long-duration geo-magnetic storms, with sawtooth events, and may sometimes be misidentified as isolated CMEs. ",
        "introduction": "Coronal mass ejections (CMEs) are the major driver of intense geo-magnetic activity \\citep[]{Richardson:2001} and have been studied extensively for the past 40 years. CMEs are observed remotely by coronagraphs and heliospheric imagers and measured {\\it in situ} by spacecraft such as ACE and {\\it Wind}. With the launch of SOHO and STEREO, the availability of white-light imagers with wide fields-of-view has made it possible to associate eruptions observed in the corona to CMEs measured {\\it in situ} at 1~AU \\citep[see for example the list by][]{Richardson:2010}. CMEs measured {\\it in-situ} may be divided into three broad categories: magnetic clouds (MCs), non-MC isolated ejecta, and complex ejecta \\citep[similar to the categories of][]{Zurbuchen:2006}. MCs have well-defined properties \\citep[see][]{Burlaga:1981}. Non-MC isolated ejecta typically have some but not all the properties of MCs, and may be sometimes referred to as MC-like ejecta \\citep[]{Lepping:2005}. They may correspond to a distorted CME or to the crossing through the ``leg'' of a CME. Lists of MCs and MC-like ejecta measured at 1~AU by the Wind and ACE spacecraft are maintained \\citep[]{Lepping:2005,Jian:2006,Richardson:2010}.  Complex ejecta result from the interaction of successive CMEs \\citep[]{Burlaga:1987}. Some consist of many individual eruptions and it is impossible to relate {\\it in situ} measurements to coronagraphic observations of CMEs \\citep[]{Burlaga:2002}. Others are made up of two clearly distinct MCs separated by an interaction region \\citep[multiple-MC events, see][]{Wang:2003, Lugaz:2005b}. Complex ejecta tend to have long duration and may drive the magnetosphere for an extended period. \\citet{Xie:2006}, for example, studied long (3 days or more) and intense (peak Dst $\\le -100$ nT) geomagnetic storms and found that 24 out of  37 such storms were associated with multiple CMEs.  While a typical CME passes over Earth in $\\sim 20$ hours, some events have duration well in excess of 30 hours \\citep[]{Marubashi:2007}. It is possible that some of these long-duration events, believed to be associated with a single, isolated CME are in fact the results of the interaction of two CMEs, a possibility raised for the 2005 May 15 CME by \\citet{Dasso:2009}.  Here, we identify a new type of complex ejecta due to the interaction of two CMEs, which results in a long-duration event with a smooth rotation of the magnetic field vector. In section \\ref{simu}, we present the result at 1~AU of two simulations, one of an isolated CME and one of two interacting CMEs, and we discuss the expected geo-effectiveness of such events. In section \\ref{geo}, we present  measurements of the 2001 March 19--22 period, which may be associated with the interaction of two CMEs in a way similar to that of the simulation. We discuss our findings and conclude in section \\ref{conclusion}. ",
        "conclusions": "By combining numerical simulations and the analysis of {\\it in situ} measurements, we have identified a new class of complex ejecta resulting from the interaction of two CMEs. Due to different orientations of the ejections, measurements at 1~AU appear to indicate the passage of a long magnetic cloud, but are in fact, due to two successive and interacting CMEs. With an appropriate orientation of the two CMEs, such an event may result in long-duration geomagnetic storm and be associated with sawtooth event.  In details, we have presented the results at 1~AU of a simulation of the interaction of two CMEs with different orientations. We have shown that the resulting complex ejecta is very similar to a MC from an isolated CME, except for the presence of a long ``tail'' in the magnetic field and the hotter temperature throughout the ejecta. We have estimated the expected Dst index for this complex ejecta, and we have found that, while the peak Dst is not as low as that from a well-oriented isolated CME, the tail in the magnetic field results in the Dst to be below $-100$~nT for more than a day, or about 50\\% longer than for the isolated CME.  We have also presented the analysis of one long-duration magnetic ejecta observed at 1~AU in 2001 March 19--22. This event resulted in a long, intense geomagnetic storm with a peak Dst of only $-149$~nT but the Dst stayed below $-50$~nT for more than 2 days. There were also a number of sawteeth in March 20 in the first half of the ejecta. Most studies have identified this as an isolated magnetic cloud with a duration in excess of 2 days. We have presented some potential evidence that this ejecta is in fact a complex ejecta associated with two CMEs.  A more complete investigation of combined {\\it in situ} and remote-sensing database will be required to assess how common this type of complex ejecta is. This could be helped in the future by the availability of remote-sensing observations of CMEs as they propagate and interact on their way to Earth \\citep[]{CShen:2012, Lugaz:2012b}. The other event that we have tentatively identified in 2004 April 3--6 was also associated with an extended sawtooth event, although the Dst index peaked only at $-117$~nT and was below $-50$~nT for only 15 hours. Further studies are also required to determine how this type of complex ejecta affects Earth's magnetosphere and how the interaction differs from that with an isolated MC or a multiple-MC event."
    },
    "1402/1402.3593_arXiv.txt": {
        "abstract": "To shed light on the fundamental problems posed by Dark Energy and Dark Matter, a large number of experiments have been performed and combined to constrain cosmological models. We propose a novel way of quantifying the information gained by updates on the parameter constraints from a series of experiments which can either complement earlier measurements or replace them. For this purpose, we use the Kullback-Leibler divergence or relative entropy from information theory to measure differences in the posterior distributions in model parameter space from a pair of experiments. We apply this formalism to a historical series of Cosmic Microwave Background experiments ranging from Boomerang to WMAP, SPT, and Planck. Considering different combinations of these experiments, we thus estimate the information gain in units of bits and distinguish contributions from the reduction of statistical errors and the `surprise' corresponding to a significant shift of the parameters' central values. For this experiment series, we find individual relative entropy gains ranging from about 1 to 30 bits. In some cases, e.g. when comparing WMAP and Planck results, we find that the gains are dominated by the surprise rather than by improvements in statistical precision. We discuss how this technique provides a useful tool for both quantifying the constraining power of data from cosmological probes and detecting the tensions between experiments. ",
        "introduction": "% \\label{sec:introduction}  \\noindent Over recent decades, observational evidence in support of the $\\Lambda$CDM model has grown steadily. Though some of the key ingredients of the model, including Dark Matter and Dark Energy, are not fully understood, an impressive array of new experiments show findings consistent with predictions of the model. Chief among the datasets are high-precision measurements of the Cosmic Microwave Background (CMB) \\cite{Wright:1996um,MacTavish:2008tm,Bennett:2012wy,Hinshaw:2012vg,Story:2012vh,Das:2011gd}. This area has received significant attention recently with the release of the first cosmological analysis of data from the Planck satellite \\cite{Collaboration:2013ww,Collaboration:2013uv}. This experiment can be seen as the latest in a long line of measurements that have targeted the CMB. At each step, data has been used to place constraints on the parameters of the $\\Lambda$CDM model using Bayesian inference to represent the constraints as a probability density in parameter space called the posterior distribution. To judge the progress made between successive measurements, a framework for comparing probability distributions is needed.  One method to quantify the difference between the constraints from different surveys is the relative entropy or Kullback-Leibler divergence \\cite{Kullback:1951va} between the respective distributions. Initially motivated from information theory, relative entropy has been proposed in the cosmology literature for forecasting and experiment design \\cite{Paykari:2013dd,Amara:2013wv,March:2011ij,0004-637X-771-1-12} as well as for parameter estimation and model selection \\cite{Verde:2013hp,Kunz:2006de}. In this paper, the relative entropy is introduced as a new tool for measuring the information gained from individual experiments by applying it to their posteriors on the full cosmological parameter space. Two distinct cases of data combinations are analyzed: adding complementary data to existing constraints and replacing data with a more accurate but correlated measurement.  The relative entropy between two posteriors measures gains in statistical precision and shifts of confidence regions at the same time. Disentangling these contributions is of great interest for detecting tensions between datasets. In the limit of linear models and Gaussian likelihoods, it is shown that the relative entropy can indeed be separated into an expected part measuring the improvements in precision and a contribution from shifts in the distribution means that is named `surprise'. Explicit expressions for the relative entropy and its decomposition into expected relative entropy and surprise are derived in this limit and can be evaluated from moments of the posteriors.  These concepts are then applied to the posteriors of the $\\Lambda$CDM parameters from the Balloon Observations of Millimetric Extragalactic Radiation and Geophysics (BOOMERANG) \\cite{MacTavish:2008tm}, the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Bennett:2012wy,Hinshaw:2012vg}, the South Pole Telescope (SPT) \\cite{Story:2012vh}, and the Planck \\cite{Collaboration:2013ww,Collaboration:2013uv} CMB surveys. Using the Monte Carlo Markov chain framework \\verb|CosmoHammer| \\cite{Akeret:2012uk}, estimates for the relative entropy, its expected, and its surprise contributions are given for different combinations of these datasets. The concepts can be easily applied to other probes, too.  This paper is organized as follows. In section \\ref{sec:information_in_parameter_estimation} the connection between relative entropy and parameter estimation is discussed. The results for disentangling expected relative entropy and surprise in the Gaussian limit are derived in section \\ref{sec:the_gaussian_limit}. Applying those concepts to CMB surveys, numerical results for the relative entropy between BOOMERANG, WMAP, SPT, and Planck data are shown in section \\ref{sec:information_gains_from_cmb}. The conclusions are summarized in section \\ref{sec:discussion}. ",
        "conclusions": "% \\label{sec:discussion}  \\noindent In order to compare the cosmological parameter constraints from different experiments, a tool for quantifying changes in posterior distributions on the full parameter space is needed. Motivated from information theory, the concept of relative entropy measures differences between distributions in a parametrization independent way and is therefore able to quantify the information gained from new data. In this work, relative entropy is used to develop a new tool for comparing the parameter constraints of the $\\Lambda$CDM model from different CMB surveys. Two ways of combining data from different experiments are discussed: complementary datasets that can be analyzed sequentially and correlated measurements that replace earlier datasets.  Relative entropy captures both changes in confidence volumes and location of the regions of the posteriors. In the regime of Gaussian likelihoods and linear models, these contributions can even be distinguished as an expected relative entropy measuring differences in confidence volume and a surprise coming from shifts in parameter space. This Gaussian regime is furthermore at least a good approximation for CMB data analysis. The notions of expected relative entropy and surprise turn the relative entropy into a powerful diagnostic for the consistency of datasets.  The relative entropy gains in units of bits from BOOMERANG, WMAP, SPT, and Planck surveys range from about 1 to 30. In general, the numbers are driven by the contributions from the expected relative entropy, $\\ere$, but in three cases the surprise is found to dominate the results. In terms of expected relative entropy, the step from Boomerang to WMAP is the biggest ($\\ere \\sim 18$ bits), followed by the update of WMAP by Planck data ($\\ere \\sim 7$ bits). The addition of SPT data to the WMAP constraints leads to an expected relative entropy gain of $\\ere \\sim 2$ bits.  Looking at the total relative entropy gains, inclusion of Planck data shows the biggest gains ($D=29.8$ bits from WMAP 9 and $D=27.8$ bits from WMAP 9 and SPT). When these numbers are decomposed, the relative entropy is found to be dominated by the surprise ($S=21.9$ and $S=21.2$ bits, respectively). These are very significant surprise values since they are $6.5$ standard deviations from expectations. Note that the expected distribution of $D$ is non-Gaussian; for the corresponding p-values see Table \\ref{tab:relent}. This indicates that the shifts in the confidence intervals of the posteriors are large compared to the shifts expected from the increased precision. This conclusion is further supported by the changes in the marginalized posterior plots and points to possible tensions when fitting predictions from $\\Lambda$CDM to different measurements, in line with other findings \\cite{Verde:2013hp,Spergel:2013va}. Other updates considered here also show significant surprise. In particular the update from WMAP 3 to WMAP 5 shows a large surprise of $S=5.5$ which is $5.3$ standard deviations away from expectations. This result might be caused by both the deviations of the WMAP 3 posterior from a normal distribution and the adjustments of the likelihood function for the low-$\\ell$ temperature power spectrum by the WMAP team. It is also interesting to note that while it is possible and expected for the surprise to be both positive and negative, the findings presented here typically show positive surprise. The reason for this is unclear, but if measurement errors are systematically underestimated this in itself would tend to bias the results towards positive surprise.  To conclude, the relative entropy is found to be a valuable diagnostic to compare constraints from different measurements. In cases where the likelihood is close to Gaussian and the model is effectively linear, the contributions from shifts in the confidence regions can be separated from the gains in precision. The resulting quantities are easy to estimate and are capable of describing the overall changes in multidimensional constraints in an efficient way."
    },
    "1402/1402.3503.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small Protoplanetary disks composed of dust and gas are ubiquitous around young stars and are commonly recognized as nurseries of planetary systems. Their lifetime, appearance, and structure are determined by an interplay between stellar radiation, gravity, thermal pressure, magnetic field, gas viscosity, turbulence, and rotation. Molecules and dust serve as major heating and cooling agents in disks. Dust grains dominate the disk opacities, reprocess most of the stellar radiation, and shield molecules from ionizing UV/X-ray photons. Disks also dynamically evolve by building up planetary systems which drastically change their gas and dust density structures. Over the past decade significant progress has been achieved in our understanding of disk chemical composition thanks to the upgrade or advent of new millimeter/Infrared facilities (SMA, PdBI, CARMA, Herschel, {\\dbf e-VLA, ALMA).} %which provide more accurate molecular observations. Some major breakthroughs in our comprehension of the disk physics and chemistry have been done since PPV. This review will present and discuss the impact of such improvements on our understanding of the disk physical structure and chemical composition. \\\\~\\\\~\\\\~} %leave this in to get the correct vertical space after the abstract ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0251_arXiv.txt": {
        "abstract": "Neutron stars and black holes are the most compact astrophysical objects we can think of and as a consequence they are the main sources of gravitational waves.  There are many astrophysically relevant scenarios in which these objects are immersed in or endowed with strong magnetic fields, in such a way that gravitational perturbations can couple to electromagnetic ones and can potentially trigger synergistic electromagnetic signatures. In a recent paper we derived the main equations for gravito-electromagnetic perturbations and studied in detail the case of polar electromagnetic perturbations driven by axial gravitational perturbations.  In this paper we deal with the case of axial electromagnetic perturbations driven by polar black-hole or neutron stars oscillations, in which the energy emitted in case is considerably larger than in the previous case. In the case of neutron stars the phenomenon lasts considerably longer since the fluid acts as an energy reservoir that shakes the magnetic field for a timescale of the order of secs. \\PACS{04.30.-w, 04.40.Nr, 95.85.Sz} ",
        "introduction": "\\label{sec:Intro} Gravitational wave observations will enable an unprecedented view of the Cosmos thanks to interferometric detectors such as the forthcoming second-generation ground-based detectors LIGO, Virgo, and KAGRA~\\cite{LIGO,VIRGO,KAGRA}, third-generation future ones like the Einstein Telescope~\\cite{ET}, and space-based ones like eLISA and DECIGO~\\cite{eLISA,DECIGO}.  In addition, there are also good prospect of detection of gravitational waves using several Pulsar Timing Arrays that are organized in the International Pulsar Timing Array collaboration~\\cite{IPTA}.  When the second generation of ground-based detectors becomes fully operational and reaches the designed sensitivity, these instruments will be truly tools of astronomical discovery. For compact objects (neutron stars and black holes), we will be able to determine from gravitational-wave observations properties such as their masses, radius, rotation rates, and, for the case of neutron stars, the structure of matter itself (i.e. equation of state)~\\cite{AK1998,KAA2001,Sotani2004,Erich2011,Sotani2011a}. Another remarkable capability of gravitational-wave observations is their potential to validate whether General Relativity is indeed the correct classical theory of gravity and to put constraints to alternative theories~\\cite{SK2004,S2009,YYT2012}. In summary, the detection and characterization of gravitational waves will provide insights on a plethora of physical phenomena, insights that will complement our knowledge obtained from particle accelerators and electromagnetic observations.  Processes in compact objects associated with strong magnetic fields are an example of physical phenomena for which only multi-messenger observations provide a complete picture. For instance, in the case of the giant flares in soft gamma-ray repeaters, evidence points to magnetars (strongly magnetized neutron stars) as the objects where these events takes place. % Moreover, in the three giant flares detected so far, theoretical studies~\\cite{L2007,SKS2007,CBK2009,CSF2009,AGS2009,S2011,HoLevin11,GCFES2012,CK2012,SNIO2012a,SNIO2012b,HoLevin12} suggest that the quasi-periodic oscillations frequencies, measured during the afterglow of the flare activity~\\cite{WS2006}, are associated with crustal oscillations in the neutron star and/or oscillations in the magnetic field of the neutron star. Furthermore, recent  studies~\\cite{A2011} have speculated that the flare activity could couple with the dynamical gravitational fields of the star to trigger the production of gravitational waves. In addition, the importance of electromagnetic counterparts to gravitational-wave signals has been also discussed in the context of recent numerical simulations of binary black hole systems~\\cite{PL2010,PT2010}, of collapsing hypermassive neutron stars~\\cite{LP2012}, and of binary neutron star systems~\\cite{PL2013a,PL2013b}.   Further studies are needed  to investigate the conditions under which electromagnetic and gravitational perturbations couple to generate emission of radiation through both channels. We have taken a first step in recent work~\\cite{SKLS2013} (Paper I). We studied the coupling of axial gravitational and electromagnetic perturbations for both, black holes and neutron stars, neglecting the back reaction of the electromagnetic waves since the energy budget in the system is dominated by gravity. Our study found that the emitted energy in electromagnetic waves driven by the gravitational waves is proportional to not only the emitted energy in gravitational waves but also to the square of the strength of the magnetic field of the central object. That is, $E_{\\rm EM}/E_{\\rm GW} = \\alpha (B/10^{15} {\\rm G})^2$, where $B$ is the magnetic field strength. For the case of a black hole background it was found that $\\alpha \\approx 8\\times 10^{-6}$, while for neutron stars this factor depends on the compactness of the neutron star.  For a compactness of $M/R = 0.162$ we obtained $\\alpha \\approx 1.61\\times 10^{-5}$ while for a compactness of $M/R = 0.237$ the result was $\\alpha \\approx 4.37\\times 10^{-6}$. Additionally, it is suggested that there can be resonances between the two types of waves under certain conditions \\cite{KT2013}.   In this paper, we investigate electromagnetic emission driven by polar gravitational perturbations. Unlike axial gravitational perturbations, the expectation is that polar perturbations are able to drive longer emission in both gravitational and electromagnetic waves. Here again, we ignore back-reaction effects from the magnetic field on the spacetime geometry.  Following the previous study in Paper I, we consider both the cases of black holes and neutron stars.  We examine in particular the dependence of the gravitational-electromagnetic coupling for neutron stars on the stellar mass and the equation of state. Then, we find that $\\alpha=1.78\\times 10^{-5}$ for a black hole background, a value roughly twice the one found in Paper I for electromagnetic waves driven by the axial gravitational waves. On the other hand, we also find that $\\alpha$, for a neutron star background, can be written as a function of ${\\cal \\chi} \\equiv (1-2M/R)^{-1}(M/R^3)^{-1}$ as $\\alpha \\times 10^7 = 0.922 - 0.903({\\cal \\chi}/{500}) + 0.552 ({\\cal \\chi}/{500})^2$, where $M$ an $R$ denote the stellar mass and radius.    This paper is organized as follows. In the next section (Section~\\ref{sec:Background}), we briefly describe the equilibrium configurations that provide the gravitational background (spacetime background geometry) as well as the the perturbative equations. In Section~\\ref{sec:Results}, we show and discuss the numerical results for both the black hole and neutron star cases. We make conclusions in Section~\\ref{sec:conclusion}. Throughout this paper we adopt geometric units, $c=G=1$, where $c$ and $G$ denote the speed of light and the gravitational constant, respectively, and use the metric signature $(-,+,+,+)$. ",
        "conclusions": "\\label{sec:conclusion}  This paper presents results of the emission of gravitational and electromagnetic waves from the strongly magnetized compact objects, extending the study presented in Paper I~\\cite{SKLS2013}. In Paper I we formulated the problem and examined in detail the case of ``electric-type\" (polar) electromagnetic waves driven by ``magnetic-type'' (axial) gravitational perturbations. In the present work, we focused on the case in which ``magnetic-type\" (axial) electromagnetic waves are driven by ``electric-type'' (polar) gravitational waves. As we have pointed out before, for the case of neutron stars, this situation is observationally more interesting because such type of gravitational waves are associated with fluid oscillations, and thus they will tell us about properties of the structure of the star.  In the case of black holes, we find that the emission of ``magnetic-type\" (axial) electromagnetic waves driven by ``electric-type'' (polar) gravitational perturbations is very similar to that of ``electric-type\" (polar) electromagnetic waves driven by ``magnetic-type'' (axial) gravitational perturbations studied in Paper I. More specifically, we obtained the same relationship as in Paper I between the radiated energy in electromagnetic waves and the energy emitted in gravitational waves, namely $E_{\\rm EM} = \\alpha B_{15}^2E_{\\rm GW}$. On the other hand, for neutron stars, we observe that the polar gravitational waves last longer than in the axial case, and thus also the radiation of  ``magnetic-type\" (axial) electromagnetic waves. This is due to the coupling with the fluid which acts as a storage reservoir of energy which slowly leaks energy to the electromagnetic sector during a longer time period.   This process is revealed via the study of the spectra of the perturbations, where we can see that the oscillation frequencies of gravitational and electromagnetic waves are the same as those of the fluid perturbations. Our results also show that the larger the compactness of the star, the stronger is the coupling between the two types of perturbations leading to an enhancement in the energy transfer and the emission of electromagnetic waves.  At the same time, the emitted radiation get damped out faster for large compactness models than for those with smaller values. Unfortunately, at frequencies of a few kHz, the emitted radiation in electromagnetic waves is easily absorbed by the interstellar medium and/or the plasma around the compact objects. To be detectable, one would have to explore reprocessing mechanisms.  We are currently investigating situations/mechanism where this electromagnetic emission may be detectable and the type of information that we could extract from them."
    },
    "1402/1402.6850_arXiv.txt": {
        "abstract": "{Towards the high galactic latitude sky, the far-infrared (FIR) intensity is tightly correlated to the total hydrogen column density which is made up of atomic ($\\ion{H}{i}$) and molecular hydrogen (H$_{2})$. Above a certain column density threshold, atomic hydrogen turns molecular.} {We analyse gas and dust properties of intermediate-velocity clouds (IVCs) in the lower galactic halo to explore their transition from the atomic to the molecular phase. Driven by observations, we investigate the physical processes that transform a purely atomic IVC into a molecular one.} {Data from the Effelsberg-Bonn $\\ion{H}{i}$-Survey (EBHIS) are correlated to FIR wavebands of the Planck satellite and IRIS. Modified black-body emission spectra are fitted to deduce dust optical depths and grain temperatures. We remove the contribution of atomic hydrogen to the FIR intensity to estimate molecular hydrogen column densities.} {Two IVCs show different FIR properties, despite their similarity in $\\ion{H}{i}$, such as narrow spectral lines and large column densities. One FIR bright IVC is associated with H$_{2}$, confirmed by $^{12}$CO $(1\\rightarrow0)$ emission; the other IVC is FIR dim and shows no FIR excess, which indicates the absence of molecular hydrogen.} {We propose that the FIR dim and bright IVCs probe the transition between the atomic and molecular gas phase. Triggered by dynamical processes, this transition happens during the descent of IVCs onto the galactic disk. The most natural driver is ram pressure exerted onto the cloud by the increasing halo density. Because of the enhanced pressure, the formation timescale of H$_{2}$ is reduced, allowing the formation of large amounts of H$_{2}$ within a few Myr.} ",
        "introduction": "\\label{sec:introduction}  The Infrared Astronomical Satellite (IRAS) showed that, up to a certain threshold, the $\\ion{H}{i}$ $21\\,\\mathrm{cm}$ emission correlates linearly with the far-infrared (FIR) dust continuum at high galactic latitudes \\citep{Low1984,Boulanger1988,Boulanger1996}. Despite this linear relationship, previous studies towards high galactic latitudes \\citep[e.g.][]{Desert1988,Reach1998} find excess FIR radiation associated with molecular gas, which is traced by carbon monoxide (CO) emission \\citep{Magnani1985}. Several of these objects show radial velocities relative to the local standard of rest (LSR) between $30\\,\\mathrm{km}\\,\\mathrm{s}^{-1} \\leq |v_{\\mathrm{LSR}}| \\leq 90\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ that are difficult to account for with a simple model of galactic rotation. These clouds are classified as intermediate-velocity clouds (IVCs) and, as a subcategory, intermediate-velocity molecular clouds (IVMCs).  Intermediate-velocity clouds are thought to originate from a galactic fountain process \\citep{Bregman2004}. Typically they show up with metal abundances close to solar, a traceable dust content, and distances below $5\\,\\mathrm{kpc}$ \\citep{Wakker2001}. Commonly, IVCs emit in the FIR \\citep[e.g.][]{Planckcollaboration2011XXIV}. Towards many IVCs a diffuse H$_{2}$ column density of $N_{\\mathrm{H}_{2}} = 10^{14}-10^{16}\\,\\mathrm{cm}^{-2}$ is inferred by interstellar absorption line measurements \\citep{Richter2003,Wakker2006}; however, only a few IVMCs are known \\citep{Magnani2010}. Dynamical processes resulting from the motion of IVCs through the halo and their descent onto the galactic disk are thought to play a major role in the process of H$_{2}$ formation in IVCs \\citep{Odenwald1987,Desert1990a,Gillmon2006b,Guillard2009}. The transition between the atomic and the molecular phase happens at a total hydrogen column density of $N_{\\mathrm{H}} = 2.0-5.0 \\times 10^{20}\\,\\mathrm{cm}^{-2}$ \\citep{ Savage1977,Reach1994,Lagache1998,Gillmon2006a,Gillmon2006b,Planckcollaboration2011XXIV}.  In this paper we study the gas and dust properties of IVCs and their transition from the atomic to the molecular phase by correlating the $\\ion{H}{i}$ distribution to the FIR dust emission at high galactic latitudes. We are interested in the conditions required for the transition from $\\ion{H}{i}$ to H$_{2}$ in IVCs at the disk-halo interface.  We use data from the new Effelsberg-Bonn $\\ion{H}{i}$-Survey \\citep[EBHIS,][]{Winkel2010,Kerp2011}, the Planck Satellite \\citep{Planckcollaboration2013I}, and IRIS \\citep{MivilleDeschenes2005b}. Our field of interest has a size of $25 \\times 25\\,\\mathrm{deg}^{2}$ centred on galactic coordinates ($l$, $b$) $=$ ($235^{\\circ}$, $65^{\\circ}$). This area is located between the Intermediate-Velocity (IV) Arch and IV Spur which are two large structures in the distribution of IVCs \\citep{Wakker2004}. In \\citet{Planckcollaboration2013XXX} this field is used also in the analysis of the cosmic infrared background.  Section \\ref{sec:data} gives the properties of our data in $\\ion{H}{i}$ and the FIR. Section \\ref{sec:methods} describes the $\\ion{H}{i}$-FIR correlation in more detail. Section \\ref{sec:entire-field} presents the observational results inferred from the $\\ion{H}{i}$, the FIR, and their correlation for the entire field, while Sect. \\ref{sec:individual-clouds} gives results for two particular IVCs within the field. Section \\ref{sec:discussion} discusses a dynamically driven $\\ion{H}{i}$-H$_{2}$ transition with respect to the two IVCs and Sect. \\ref{sec:conclusion} summarises our results. ",
        "conclusions": "\\label{sec:conclusion}  We correlate $\\ion{H}{i}$ emission to the brightness at various wavelengths in the FIR dust continuum using new data from the  Effelsberg-Bonn $\\ion{H}{i}$ Survey (EBHIS) and the Planck satellite complemented by IRIS data. We study in detail two IVCs that show many similarities in their $\\ion{H}{i}$ properties, such as narrow spectral lines with $\\mathrm{FWHM}\\simeq5\\,\\mathrm{km}\\,\\mathrm{\\mathrm{s}}^{-1}$ and $\\ion{H}{i}$ column densities of $N_{\\ion{H}{i}} = 3 - 4 \\times 10^{20}\\,\\mathrm{cm}^{-2}$. Despite their similarity in $\\ion{H}{i}$, their FIR emission exhibits large differences: one cloud is FIR bright while the other IVC is FIR faint.  From the quantitative correlation of $\\ion{H}{i}$ and FIR emission, we calculate maps of molecular hydrogen column density revealing large amounts of H$_{2}$ in the field of interest for which no existing CO surveys of the region has detected a CO counterpart. How much of this H$_{2}$ is actually CO-dark gas we cannot tell since the CO survey data that is available today is not sensitive enough. We do, however, know that the molecular IVC contains CO. The $\\ion{H}{i}$ emission traces only a part of the total gas distribution. Together with the inferred H$_{2}$ column densities, the relation between gas and dust is consistent.  Based on our findings we describe a scenario of a dynamical transition from atomic to molecular IVCs in the lower galactic halo. During the descent of the IVCs through the galactic halo, they are compressed as a result of external pressure and ram pressure. Once the ram pressure exceeds the thermal pressure, shocks are created which enhance the pressure locally and accumulate gas and dust. An increased pressure reduces the formation timescale of H$_{2}$ in condensations of the cold atomic medium.  The only physical distinctions between the two IVCs are the amount of dust within the clouds, measured by the dust surface mass density, and the amount of H$_{2}$. The molecular IVC has a factor of $2-3$ more dust within its central region than the atomic IVC, which is accompanied by a $50\\%$ higher total hydrogen column density. On the other hand, IVC\\,1 has more $\\ion{H}{i}$ mass in total. Apparently, this mass is not distributed so as to allow efficient H$_{2}$ formation. According to the data presented in this paper, we expect that the atomic IVC will also turn molecular in a few Myrs.  Processes on spatial scales that are not resolved by our data govern the evolution of an IVC from an atomic to a molecular cloud. The $10.8\\,\\arcmin$ resolution of EBHIS corresponds to a spatial resolution of $1\\,\\mathrm{pc}$ at a distance of $0.4\\,\\mathrm{kpc}$. However, the accumulation and condensation of smaller and denser clumps regulate the H$_{2}$ formation on sub-parsec scales. Radio-interferometric observations should reveal a different distribution of $\\ion{H}{i}$ in the two IVCs on sub-parsec scales, for example compact cores in the molecular IVC.  Our approach appears to open a way to search for dark H$_{2}$ gas across the entire sky. Globally, this new search may reveal other clouds in transition from the atomic to the molecular gas phase. Compared to the low angular resolution of former large-scale $\\ion{H}{i}$ single-dish surveys, the few detections of molecular IVCs so far may be due to the small-angular extent of the molecular cores."
    },
    "1402/1402.4408_arXiv.txt": {
        "abstract": "Atmospheric spectroscopy of extrasolar planets is an intricate business. Atmospheric signatures typically require a photometric precision of $1 \\times 10^{-4}$ in flux over several hours. Such precision demands high instrument stability as well as an understanding of stellar variability and an optimal data reduction and removal of systematic noise. In the context of the \\echo~mission concept, we here discuss the  data reduction and analysis pipeline developed for the \\echo~end-to-end simulator \\echosim. We present and discuss the step by step procedures required in order to obtain the final exoplanetary spectrum from the \\echosim~`raw data' using a simulated observation of the secondary eclipse of the hot-Neptune 55 Cnc e. ",
        "introduction": "\\label{intro}  The field of extrasolar planets is innovative as it is new. Recent successes in characterisation of extrasolar planets are also always tales of characterising the instrument response function to an unprecedented detail. Always being at the edge of technical feasibility means that instrument calibration, observing strategy as well as data analysis and modelling are interdependent.  In the light of the \\echo~ESA-M3 mission concept \\cite{tinetti12}, such interdependence becomes important in the study of engineering decisions and instrument trade-offs. In other words, one needs to simulate the full observational and data analysis chain in order to gauge the impact the instrument concept has on the achievable error bar of the detection. Such a feat requires an advanced mission end-to-end simulator as well as an advanced data analysis pipeline. In this paper, we discuss the data analysis pipeline which is used in conjunction to the  mission simulator, \\echosim~ \\cite{pascale13}. The \\echosim~data pipeline (from here on \\echodp) is a stand-alone software custom built for \\echosim~but with easy adaptability to other instruments and data sets in mind.  The method by which the \\echo~mission will characterise the nature of extrasolar planets is by time resolved spectroscopy of their atmospheres, in particular of transiting extrasolar planets. Briefly, when an exoplanet transits in front of its host star (in our line of sight) we observe a diminishing of the stellar flux due to the obscuration of the planet. The depth of the resulting lightcurve allows us to estimate the planetary radius (given the stellar radius is known). This we refer to as transit (or primary eclipse) observation. Should the exoplanet feature an extended atmosphere, we expect some of the stellar light to filter through the terminator region of the planetary atmosphere. Here we are sensitive to molecules absorbing the stellar light at specific wavelengths. We hence perceive a variation of transit depths depending on the wavelength range observed. These variations constitute the signatures of an exoplanetary absorption spectrum. Similarly, we can observe the occultation (or secondary eclipse) where the thermal contribution of the exoplanet's day-side is lost to the observer as the planet passes behind its host star. The study of transmission and emission spectroscopy is now a well established field for both space and ground based observations of exoplanetary atmospheres (e.g. \\cite{beaulieu10,beaulieu11,charbonneau08,brogi12,bean11,swain08,swain08b,swain09a,crouzet12,deming13,grillmair08,thatte10,tinetti07,pont08,swain12,knutson11,sing11,tinetti10,mooij12,bean11b,stevenson10} also see \\cite{tinetti13} for a comprehensive review).  \\subsection{\\echosim}  \\echosim~ is the \\echo~ mission end-to-end simulator. \\echosim~  implements a detailed simulation of the major observational and instrumental effects, and associated systematics. It also allows the influence of individual instrumental and astrophysical parameters to be studied and thus represents a key tool in the optimisation of the instrument design. Observation and calibration strategies, data reduction pipelines and analysis tools can all be designed effectively using the realistic outputs produced by \\echosim \\cite{pascale13,waldmann13}. The simulation output closely mimics standard STSci\\footnote{$http://archive.stsci.edu/hst/$} FITS files, allowing for a high degree of compatibility with standard astronomical data reduction routines.  \\subsection{Examples}  We illustrate individual steps in \\echodp~using diagrams. Unless specified otherwise, we follow a single data processing run of \\echosim~simulated data of the hot-Neptune 55 Cnc e. \\echosim~was run to simulate the Chemical Census mode of \\echo, in which we co-add (in the case of 55 Cnc e) five eclipse observations to obtain a minimal signal-to-noise  (S/N) of the final spectrum of S/N $\\sim$ 5.  For this we assume spectra reconstructed with a resolving powers of 50,50,30,30,30 for the VNIR, SWIR, MWIR-1, MWIR-2, and LWIR channels. With each channel having a larger native resolving power, this allows us to increase the SNR of the detection for this particular observing mode. See \\cite{tinetti12} for a review of the proposed \\echo~observing modes. ",
        "conclusions": "\\echodp~is a custom built data reduction and analysis pipeline for the \\echosim~end-to-end mission simulator of the \\echo~mission concept.  Despite its customised nature, we have developed the pipeline with easy adaptability (through its fully object-orientated programming ) to other instruments and data-sets in mind. The pipeline features state of the art data de-correlation algorithms as well as a full Bayesian analysis implementation via adaptive MCMC. Both these aspects, the de-trending as well as the exploration of stellar variability are not required for the current version of \\echosim~(version 3.x) but included with future releases. These releases will have special emphasis on realistic stellar noise simulations \\cite{ballerini12} as well as more advanced non-Gaussian instrument systematics."
    },
    "1402/1402.1707_arXiv.txt": {
        "abstract": "Atom interferometry is an exciting tool to probe fundamental physics. It is considered especially apt to test the universality of free fall by using two different sorts of atoms. The increasing sensitivity required for this kind of experiment sets severe requirements on its environments, instrument control, and systematic effects. This can partially be mitigated by going to space as was proposed, for example, in the Spacetime Explorer and Quantum Equivalence Principle Space Test (STE-QUEST) mission. However, the requirements on the instrument are still very challenging. For example, the specifications of the STE-QUEST mission imply that the Feshbach coils of the atom interferometer are allowed to change their radius only by about $260\\,$nm or $2.6\\cdot 10^{-4}\\,$\\% due to thermal expansion although they consume an average power of $22\\,$W. Also Earth's magnetic field has to be suppressed by a factor of $10^5$. We show in this article that with the right design such thermal and magnetic requirements can indeed be met and that these are not an impediment for the exciting physics possible with atom interferometers in space. ",
        "introduction": "One of the biggest challenges in current theoretical physics is that of finding a valid theory of quantum gravity. Although many theories were proposed, ultimately this question has to be resolved experimentally. Thus, many experiments regarding tests of the fundamental assumptions of gravity and its basic effects were carried out and suggested including tests of the Universality of Free Fall (UFF). This basic principle is a cornerstone of Einstein's theory of gravity and implies that the trajectories of freely falling, structureless test particles in a gravitational field only depend on their initial position and velocity. In particular, the path of such test particles is independent of their composition. A violation of this principle measured by the E\\\"otv\\\"os ratio would help to discriminate between the different proposed theories of quantum gravity and it would point in the right direction for further theoretical and experimental development. If the UFF holds to the tested accuracy, bounds will be placed on the viable alternative theories.  A very thriving tool to test the UFF using quantum matter is dual species atom interferometry,\\cite{Peters_2001,Fray_2004,Bonnin_2013} where two atomic clouds propagate freely and their trajectories are compared. Although the advances along this avenue are tremendous, very stringent requirements have to be satisfied to reach down to an accuracy of the UFF test comparable to the classical tests like lunar laser ranging\\cite{Williams2012CQG} and torsion balance experiments.\\cite{Schlamminger08PRL} In the present paper, we show that the requirements regarding the magnetic field and the thermal requirements, which are closely tied to each other, can indeed be satisfied. As a source for specific requirements we refer to the Spacetime Explorer and Quantum Equivalence Principle Space Test (STE-QUEST) mission, which is a medium-size candidate mission in the Cosmic Vision Program of the European Space Agency. Different aspects of the STE-QUEST mission are described in Ref.\\ \\onlinecite{Hechenblaikner_2013,Aguilera_2013,Tino_2013,Schuldt_2014}. The specific requirements for STE-QUEST are derived in Ref.\\ \\onlinecite{Schubert_2013}.  The STE-QUEST atom interferometer was designed with heritage from projects such as the DLR project QUANTUS\\cite{Zoest_2010} (QUANTengase Unter Schwere\\-losig\\-keit) and the CNES project I.C.E.\\cite{Nyman_2006} (Inter\\-f\\'erom\\'etrie Coh\\'erente pour l'Espace). In QUANTUS, experiments with Rubidium (Rb) Bose-Einstein condensates are carried out in the drop tower at ZARM in Bremen and are currently prepared for a sounding rocket by the end of 2014.\\cite{MAIUS} A dual species interferometer is already included in I.C.E., where the experiment is performed in parabolic zero-g flight.  The experimental sequence foreseen for STE-QUEST, is to cool two ensembles of \\textsuperscript{85}Rb and \\textsuperscript{87}Rb atoms down to Bose-Einstein condensation in two steps. First atoms are loaded from a magneto-optical trap into a magnetic trap on an atom chip for evaporative pre-cooling. Then, they are loaded into an optical dipole trap for final evaporative cooling. During this step a strong homogeneous magnetic field of $160\\,$G is applied in order to tune atomic interactions using a so called Feshbach resonance. The magnetic field is generated by two coils in Helmholtz configuration with a middle distance of $105\\,$mm, a coil's mean diameter of $206\\,$mm, and equal electrical current of $4.8\\,$A flowing in the same direction. This way, two Bose-Einstein condensates of $10^6$ atoms shall be generated and released from the trap to fall freely in Earth's gravitational field. A detailed description of the atom interferometer design of STE-QUEST can be found in Ref.\\ \\onlinecite{Schuldt_2014}.  The article is organized as follows: In Sec.\\ \\ref{sec:requirements}, the thermal requirements and the requirements for the magnetic shielding are summarized for the example of the STE-QUEST atom interferometer. Subsequently, the methods and models used for the analysis of the thermal control system and the magnetic shielding are discussed in detail. The results are summarized in Sec.\\ \\ref{sec:results}. ",
        "conclusions": "While dual species atom interferometry in space is a powerful and promising tool for fundamental physics science, the technological challenges that have to be overcome in order to realize a sensor offering a high measurement accuracy are considerable. Among these, the thermal management and the magnetic shielding of the sensor hardware have a strong impact on the achievable sensor accuracy. In this paper, we have analyzed the current design of the thermal control system and the magnetic shielding with respect to the STE-QUEST mission specification as a test case scenario. Both systems are designed and verified by means of complex 3D FE models, which have been used to simulate the evolution of heat flows, temperatures and magnetic field gradients for the dynamic boundary condition throughout the STE-QUEST orbit and the measurement cycle. The numeric results show that the current design of the atom interferometer allows to satisfy thermal and magnetic requirements by a thorough design of both the thermal control system and the magnetic shielding. Furthermore a more advanced design of the magnetic coils could help to relax the requirements on thermal expansion. The design and the FE models of both the thermal control system and the magnetic shielding will be used for future optimization or verification of overall design changes arising from other subsystem or mission requirements."
    },
    "1402/1402.5312_arXiv.txt": {
        "abstract": "{ Most Class~II sources (of nearby star forming regions) are surrounded by disks with weak millimeter continuum emission. These ``faint'' disks may hold clues to the disk dissipation mechanism. However, the physical properties of proto-planetary disks have been directly constrained by imaging only the brightest sources. } { We attempt to determine the characteristics of such faint disks around classical T Tauri stars, and to explore the link between disk faintness and the proposed disk dispersal mechanisms (accretion, viscous spreading, photo-evaporation, planetary system  formation). } { We performed high-angular resolution ($0.3''$) imaging of a small sample of disks (9 sources) with low 1.3\\, mm continuum flux (mostly $<30$\\,mJy) with the IRAM Plateau de Bure interferometer and simultaneously searched for $^{13}$CO (or CO) J=2-1 line emission. Using a simple parametric disk model, we determine characteristic sizes of the disks, in dust and gas, and we constrain surface densities in the central 50\\,AU. } { All disks are much smaller than the bright disks imaged so far, both in continuum and $^{13}$CO lines (5 detections). In continuum, half of the disks are very small, with characteristic radii less than 10\\,AU, but still have high surface density values. Small sizes appear to be the main cause for the low disk luminosity. Direct evidence for grain growth is found for the three disks % that are sufficiently resolved. Low continuum opacity is attested in two systems only, but we cannot firmly distinguish between a low gas surface density and a lower dust emissivity resulting from grain growth. Finally, we report a tentative discovery of a $\\sim 20$\\,AU radius cavity in DS Tau, which with the (unresolved) ``transition'' disk of CX~Tau, brings the proportion of ``transitional'' disks to a similar value to that of brighter sources. The existence of cavities cannot explain by itself their observed low mm flux. } { This study highlights a category of very compact dust disks, still exhibiting high surface densities, which may represent up to 25\\% of the whole disk population. While its origin is unclear with the current data alone, it may be related to the compact planetary systems found by the Kepler mission.} ",
        "introduction": "The many known exoplanet systems display tremendous variety in the masses of their stars, the masses of planets they host, and their architectures.  Study of the early evolution of protostellar disks will contribute to our understanding of the origins of this diversity. \"Faint\" disks (which we defined as having 1.3 mm continuum flux $\\simless 60$ mJy at the Taurus distance, 140 pc) play a key role in this respect. First of all, most disks around Class II sources are ``faint''. \\citet{Andre+Montmerle_1994} find that 50 \\% of the Class II sources in $\\rho$ Oph (at 120 pc) have 1.3\\,mm flux $< 40$ mJy, while only 25 \\% have flux $> 125$~mJy. Similar statistics were obtained in the Taurus region by \\citet{Beckwith+etal_1990}. Second, young stars of mass  $< 0.3$\\,M$_\\odot$ (the most numerous in our Galaxy) are expected to be surrounded by low-mass, intrinsically faint disks from theoretical considerations. Until the work of \\citet{Andrews+etal_2013}, this trend was not clearly observed because of incompleteness of the samples towards low masses (sensitivity). The large dispersion of disk continuum brightness as a function of stellar mass also loosen the expected correlation. Finally, low-mass stars with faint disks play a critical role in validating stellar evolution models. Providing accurate stellar masses in the 0.2-0.4\\,$\\Msun$ range would allow critical tests of early stellar evolution models, hopefully offering better age estimates for all young stars to ultimately establish a more reliable clock for planetary system formation.  Disks can be ``faint'' for many different reasons. \\textit{Low disk surface density} is the simplest reason. Temperature also affects continuum flux, but varies over a much more limited range. \\textit{Dust settling} when seen at high inclination  \\citep[$> 80^\\circ$, e.g. the edge-on objects HK Tau B, HV Tau C in][]{Guilloteau+etal_2011} leads to smaller flux \\citep[see the study of][]{Boehler+etal_2013}, as the warmer parts are hidden by opacity. Disks may also \\textit{just be small}, as for example \\object{BP Tau} \\citep{Dutrey+etal_2003}, sometimes as a result of \\textit{outer disk truncation}, as happens in binaries \\citep[e.g. Haro 6-10, UY Aur,][]{Guilloteau+etal_2011,Duvert+etal_2000}.  On another hand, \\textit{inner disk clearing}, which suppresses the densest parts of the disk, is also observed in many objects \\citep[e.g. AB Aur, LkCa 15, MWC 758,][]{Pietu+etal_2005,Pietu+etal_2006,Isella+etal_2010a}. Finally, \\textit{grain growth}, which reduces the opacity per unit mass, can lead to small (sub)mm continuum flux densities.  These causes are not mutually exclusive: \\object{HH 30} \\citep{Guilloteau+etal_2008} combines small size, inner hole and grain growth, as well as a somewhat lower than average surface density. The various possible causes have different observational consequences.  \\citet{Ricci+etal_2010a} found a correlation between the spectral index $\\alpha$ ($S(\\nu) \\propto \\nu^\\alpha)$ with S(1.3\\,mm), with $\\alpha \\simeq 2$ for low flux densities. $\\alpha = 2$ indicates either optically thick emission, or large grains (``pebbles''), as the dust emissivity index $\\beta = \\alpha-2$ in the optically thin case. However, the observed correlation fails to be reproduced by current models of disk evolution including grain growth and viscous evolution \\citep[][their Fig.2]{Birnstiel+etal_2010}. A trend for smaller sizes at low flux \\citep[][their Fig.7]{Andrews+etal_2010} supports the compact disk interpretation. Furthermore, since \\citet{Guilloteau+etal_2011} showed that larger grains are preferentially found in the inner 60 - 100 AU of disks, small disks are also expected to have lower $\\alpha$, producing the observed correlation between flux and spectral index.   Whatever the cause of the faintness, ``faint'' disks are essential objects for understanding the planetary system formation process. Grain growth and dust settling are major steps in this respect. Inner cavities can be signposts of planetary system formation, while disks with low surface density may be in the dust (and gas) dissipation stage. To separate between the possible causes, resolved images are essential. Because of the resolution and sensitivity limitations of the existing mm arrays, most studies have focussed on the brighter objects \\citep[e.g.][]{Andrews+Williams_2007,Isella+etal_2009,Guilloteau+etal_2011}, so that resolved images of faint disks are rare. A first attempt was performed by \\citet{Andrews+etal_2010}, who studied the continuum emission from some disks in the $\\rho$ Oph regions,  6 of which would qualify as ``faint'', showing that on average these disks appeared indeed smaller than brighter ones. However, the sensitivity was insufficient to study the CO line emission.  We report here on a high angular resolution (0.4-0.6$''$) study in line and continuum emission of 9 disks in the Taurus complex, 8 of them having $S_\\nu < 30$ mJy at 1.3\\,mm, and the last one, CW Tau, being at the boundary between ``faint'' and ``bright'' disks with $S_\\nu \\approx 60$ mJy. ",
        "conclusions": "We report the first high angular resolution study of disks with low level of continuum emission at mm wavelengths. We find that these disks are substantially smaller than the brighter ones imaged so far. A majority of them remain essentially unresolved at 0.4$''$ resolution, leading to radii $< 15$ AU. These disks are small enough to be optically thick at 1.4 mm, which provides a simple explanation for their apparent spectral index $\\alpha \\sim 2$.  $^{13}$CO emission has been detected in five disks. Although the gas distribution extends further out than the apparent size of the dust emission (as in all other disks around T Tauri and HAe stars), its outer radius is also small compared to previously studied gas disks.  Only two disks (V836 Tau, an intermediate case between Class II and Class III sources, and HO Tau) have apparently low surface densities simultaneously attested by their low opacity in continuum and by their low molecular content. These stars may be caught in the act of dissipating their disks. Despite its higher $^{13}$CO content, DS Tau, which forms a common proper motion pair with a Class III/wTT star, is another good candidate for a dissipating disk because of the possible existence of an inner cavity.  This study reveals a population of small disks whose surface density remains large in the central 10 AU region. This population had been ignored so far because of the limited sensitivity and angular resolution of previous observations. It may represent up to 25 \\% of the whole disk population in Taurus. These disks apparently only differ from their larger siblings by their small sizes. The origin for such small sizes remains unknown, and requires higher resolution observations with ALMA and the JVLA to be unveiled. Such observations would also reveal the ability of these disks to form planets."
    },
    "1402/1402.7317_arXiv.txt": {
        "abstract": "Recent studies suggest that binary neutron star (NS-NS) mergers robustly produce the heavy $r$-process nuclei above the atomic mass number $A \\sim 130$ because of their ejecta consisting of almost pure neutrons (electron fraction of $Y_\\mathrm{e} < 0.1$). However, little production of the lighter $r$-process nuclei ($A \\approx 90$--120) conflicts with the spectroscopic results of $r$-process-enhanced Galactic halo stars. We present, for the first time, the result of nucleosynthesis calculations based on the fully general-relativistic simulation of a NS-NS merger with approximate neutrino transport. It is found that the bulk of the dynamical ejecta are appreciably shock-heated and neutrino-processed, resulting in a wide range of $Y_\\mathrm{e}$ ($\\approx 0.09$--0.45). The mass-averaged abundance distribution of calculated nucleosynthesis yields is in reasonable agreement with the full-mass range ($A \\approx 90$--240) of the solar $r$-process curve. This implies, if our model is representative of such events, that the dynamical ejecta of NS-NS mergers can be the origin of the Galactic $r$-process nuclei. Our result also shows that the radioactive heating after $\\sim 1$~day from the merging, giving rise to $r$-process-powered transient emission, is dominated by the $\\beta$-decays of several species close to stability with precisely measured half-lives. This implies that the total radioactive heating rate for such an event can be well constrained within about a factor of two if the ejected material has a solar-like $r$-process pattern. ",
        "introduction": "\\label{sec:intro}  The astrophysical site of the $r$-process, the rapid neutron-capture process that makes half the elements heavier than iron, remains a long-standing mystery of nucleosynthesis. Recently, compact binary mergers (CBMs) of double neutron star (NS-NS) and black hole--neutron star (BH-NS) systems have received considerable attention as possible sources of the $r$-process nuclei \\citep{Lattimer1974, Symbalisty1982, Eichler1989, Meyer1989, Freiburghaus1999} according to the following reasons.  First, radioactively powered ``kilonova'' emission from the $r$-processed ejecta can be a promising electro-magnetic counterpart to the gravitational-wave signal from a CBM event \\citep{Li1998, Metzger2010, Goriely2011, Kasen2013, Barnes2013, Tanaka2013, Grossman2014}. The possible identification of a kilonova associated with the \\textit{Swift} GRB~130603B \\citep{Berger2013, Tanvir2013} also indicates that CBMs are the progenitors of short-duration gamma-ray bursts and the sources of $r$-process elements \\citep{Hotokezaka2013a, Tanaka2014}.  Another reason is that core-collapse supernovae (CCSNe; in particular proto-NS wind), the site that has been believed to be the promising sources of the $r$-process nuclei, are found to provide only marginal conditions for making the elements beyond iron \\citep{Martinez2012, Roberts2012, Fischer2012}. Nucleosynthesis studies with such physical environments confirm that CCSNe produce the elements only up to the atomic mass number $A \\sim 110$ \\citep{Wanajo2011, Wanajo2013}. One possible exception could be the scenario of (still hypothetical) rapidly rotating, strongly magnetized CCSN cores \\citep{Winteler2012}.  Recently, \\citet{Goriely2011} and \\citet{Bauswein2013} have explored nucleosynthesis based on the approximate (conformally flat spatial metric) general-relativistic (GR) simulations of NS-NS mergers. They found that the ejecta had extremely low electron fractions ($Y_\\mathrm{e} < 0.1$), which led to fission recycling and thus robust production of only heavy $r$-process nuclei with $A \\gtrsim 130$. Similar results were obtained from the Newtonian simulations of NS-NS and BH-NS mergers by \\citet{Roberts2011, Korobkin2012, Rosswog2014}.  Little production of the lighter $r$-process nuclei ($A \\approx 90$--120) conflicts, however, with the recent spectroscopic results of Galactic halo stars \\citep{Sneden2008, Siqueira2014}. That is, the so-called ``universality'' of the (solar-like) $r$-process pattern, first identified for $Z \\gtrsim 56$ ($A \\gtrsim 140$), persists down to $Z \\sim 38$ ($A \\sim 90$) within about a factor of two. There has been no sign of nucleosynthetic events making the nuclei exclusively with $A \\gtrsim 130$. Contribution from, e.g., the subsequent BH accretion-torus wind \\citep{Surman2008, Wanajo2012, Fernandez2013} might cure this problem.  In this Letter, we report our first result of nucleosynthesis study based on the full-GR, approximate neutrino transport simulation of a NS-NS merger. The GR effects, being crucial for the dynamical evolutions of merger ejecta as pointed out by \\citet{Hotokezaka2013b}, were not fully taken into account in the previous studies. Moreover, neutrino transport that can affect the ejecta $Y_\\mathrm{e}$ is neglected in all previous studies \\citep[except for the 2D Newtonian simulation by][without nucleosynthesis calculations]{Dessart2009}. Our NS-NS merger model is described in Section~\\ref{sec:model}. The subsequent nucleosynthesis result is presented in Section~\\ref{sec:nucleosynthesis}. The radioactive heating rates (relevant for kilonova emission) are also obtained from the nucleosynthesis calculations (Section~\\ref{sec:heating}). ",
        "conclusions": "\\label{sec:summary}  We examined $r$-process calculations based on the full-GR, approximate neutrino transport simulation of the NS-NS merger with the equal masses ($= 1.3\\, M_\\odot$) of NSs. Different from previous studies, the merger ejecta exhibited a wide range of $Y_\\mathrm{e} \\approx 0.09$--0.45 that led to the nucleosynthetic abundance distribution being in good agreement with the solar $r$-process pattern. Given that the model is representative, our result (with the present estimate of the Galactic event rate) implies that NS-NS mergers can be the major origin of all the $r$-process elements in the Galaxy.  Our result also indicates that the radioactive heating (that powers a kilonova transient) after $\\sim 1$~day from the merging is dominated by the $\\beta$-decays of a small number of species with measured half-lives. The total heating rates are thus well approximated by the $\\beta$-decays of the solar $r$-process-like abundances as well as by the approximation of $\\propto t^{-1.3}$. Detailed multi-dimensional information of nucleosynthesis abundances should be, however, taken into account when we consider the spatial dependences of kilonova emission.  Our result implies that the previous thought of NS-NS merger events, dynamically ejecting almost pure NS matter, should be reconsidered. The shock-heated and neutrino-processed ejecta from a HMNS are in fact modestly neutron-rich: the phenomenon similar to the early stage of a CCSN (a proto-NS instead of a HMNS). Much more works will be needed to test if similar results are obtained with full-3D nucleosynthetic analyses, with different NS masses and their ratios, with other (reasonable) EOSs, with higher spatial resolution, etc. Nucleosynthetic contributions from BH-NS mergers, as well as from the BH-accretion tori subsequent to NS-NS/BH-NS mergers, should be also explored to draw conclusions on the role of CBMs to the Galactic chemical evolution of the $r$-process nuclei."
    },
    "1402/1402.3759_arXiv.txt": {
        "abstract": "% High energy emissions from supernovae (SNe), originated from newly formed radioactive species, provide direct evidence of nucleosynthesis at SN explosions. However, observational difficulties in the MeV range have so far allowed the signal detected only from the extremely nearby core-collapse SN 1987A. No solid detection has been reported for thermonuclear SNe Ia, despite the importance of the direct confirmation of the formation of $^{56}$Ni, which is believed to be a key ingredient in their nature as distance indicators. In this paper, we show that the new generation hard X-ray and soft $\\gamma$-ray instruments, on board {\\em Astro-H} and {\\em NuStar}, are capable of detecting the signal, at least at a pace of once in a few years, opening up this new window for studying SN explosion and nucleosynthesis. ",
        "introduction": "Supernova (SN) explosions trigger (or are triggered by) explosive nucleosynthesis, and they are believed to be main production sites of heavy elements in the Universe. The resulting yields are sensitive to explosion mechanism(s), and thus studying nucleosynthesis products is important to uncover the still-debated explosion mechanism. Especially important is the production of $^{56}$Ni -- this is the origin of Fe (as a result of the radioactive decay chain $^{56}$Ni $\\to$ $^{56}$Co $\\to$ $^{56}$Fe), and the decay is believed to provide a source of emissions from (many classes of) SNe through thermalization of emitted $\\gamma$-rays and positrons. In type Ia supernovae (SNe Ia), about half of an exploding white dwarf in mass is processed into $^{56}$Ni, supporting their huge luminosities as distance indicators. However, the most direct evidence in this scenario is still missing -- there has been no solid detection of the decay $\\gamma$-rays from SNe except for SN 1987A (e.g., Dotani et al., 1987; Sunyaev et al., 1987). Especially, no solid detection has been reported for SNe Ia (see Milne et al., 2004 for a review).  From a theoretical point of view, studying this high energy emission has been restricted to one-dimensional models (see Milne et al., 2004, for a review) despite the importance of multi-dimensional structures of the explosion both in theory and observation (e.g., Kasen et al., 2009; Maeda et al., 2010a). Most previous studies also focused on the emission in the MeV range. In this paper, we present our radiation transfer simulations of the high energy emission based on the state-of-the-art SN Ia explosion models. We extend our analysis to hard X-ray and soft $\\gamma$-ray regimes, for which dramatic improvement is expected in the observational sensitivities thanks to new generation observatories like {\\em NuStar} (Koglin et al., 2005) or {\\em Astro-H} (Takahashi et al., 2010). We predict that these telescopes are capable of detecting the radioactive decay signals from SNe Ia, at a rate of once in a year or at least once in a few years. We also briefly comment on perspectives for core-collapse SNe. ",
        "conclusions": "\\begin{table*} [!ht] % \\begin{center} \\caption{Expected Detectability (for an exposure of $10^6$ s centered at the peak date in each band pass). Shown here are limiting distance and the expected number of SNe Ia within the distance (shown in parenthesis). `cons' and `opt' are conservative and optimistic estimates, respectively. See Maeda et al. (2012) for details.} \\bigskip \\begin{tabular}{lllll} \\hline\\hline &     &   DD2D\\_asym\\_04 & W7 & DD2D\\_iso\\_04\\\\\\hline & $M$($^{56}$Ni)/$M_{\\odot}$  & 1.02 & 0.64 & 0.42\\\\\\hline Band (keV) & Instrument & Mpc (SNe year$^{-1}$) & &\\\\\\hline 60--80 & HXI & 13.9 (0.43) & 17.7 (0.96) & 10.5 (0.09)\\\\ & NuStar (cons.) & 13.0 (0.43) & 16.5 (0.70) & 9.7 (0.09)\\\\ & NuStar (opt.) & 18.4 (1.13) & 23.3 (2.52) & 13.8 (0.43)\\\\ 158 & SPI & 4.6 ($<$0.09) & 2.9 ($<$0.09) & 2.3 ($<$0.09)\\\\ & SGD (cons.) & 22.2 (2.09) & 14.2 (0.43) & 11.4 (0.09)\\\\ & SGD (opt.) & 38.5 (6.70) & 24.6 (2.96) & 19.7 (1.57)\\\\ 200--460 & SPI & 3.7 ($<$0.09) & 2.7 ($<$0.09) & 2.3 ($<$0.09) \\\\ & SGD (cons.) & 11.6 (0.09) & 8.6 (0.09) & 7.1 (0.09) \\\\ & SGD (opt.) & 20.2 (1.74) & 14.8 (0.43) & 12.3 (0.26) \\\\ 812 & SPI & 4.3 ($<$0.09) & 2.6 ($<$0.09) & 2.0 ($<$0.09)\\\\ & GRIPS & 16.8 (0.87) & 10.0 (0.09) & 7.6 (0.09)\\\\ 847 & SPI & 7.7 (0.09) & 5.4 ($<$0.09) & 4.6 ($<$0.09)\\\\ & GRIPS & 29.8 (4.52) & 21.0 (2.00) & 18.0 (1.04)\\\\\\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*}[ht] \\centerline{ \\begin{minipage}[]{0.3\\textwidth} \\resizebox{30mm}{!}{\\includegraphics{maeda_FRAWS_2013_01_fig03a_REVISED.ps}} \\end{minipage} \\begin{minipage}[]{0.3\\textwidth} \\resizebox{30mm}{!}{\\includegraphics{maeda_FRAWS_2013_01_fig03b_REVISED.ps}} \\end{minipage} \\begin{minipage}[]{0.3\\textwidth} \\resizebox{30mm}{!}{\\includegraphics{maeda_FRAWS_2013_01_fig03c_REVISED.ps}} \\end{minipage} } \\caption{Detector response simulations for an exposure of $10^{6}$ seconds for selected models (Tab. 1), for HXI (black) and SGD (red) on board {\\em Astro-H}. The model spectra at 20 days after the explosion are used as input, placed at distances of 15 Mpc. The sensitivity curves are adopted from Kokubun et al. (2010), Tajima et al. (2010), and Takahashi et al. (2010). Note that the photon count is very low in the HXI band for all the models at this distance, thus an apparent detection by HXI (left panel) just comes from the statistical fluctuation. } \\label{fig1} \\end{figure*}   According to our simulations of radioactive decay signals from SNe Ia, the new generation hard X-ray and soft $\\gamma$-ray observatories (either {\\em NuStar} or {\\em Astro-H}) are expected to be capable of detecting these signals from SNe Ia up to $\\sim 20$ Mpc with $10^6$ s exposure. This will hopefully lead to nearly annual detections, dramatically changing the field. We thus propose follow-up of nearby SNe Ia by these telescopes in a ToO mode. With a standard set up with a few $10^{5}$ s exposure, the detection will be limited to extremely nearby objects (i.e., up to $\\sim $8 - 12 Mpc for SNe Ia with average brightness), but still there is a good chance of first solid detection of the signal from SNe Ia. Once detected, it will provide various diagnostics on explosive nucleosynthesis and explosion mechanisms, and here a combination of hard X-ray and soft $\\gamma$-ray will be essential.  A similar argument applies for core-collapse SNe. By combining results from similar simulations for core-collapse SNe (Maeda, 2006) and those obtained for SNe Ia (Maeda et al., 2012), we find the following. SNe IIp (an explosion of a red supergiant) is not a promising target in the soft bands, since the thick H envelope is still opaque in the early phase where the $^{56}$Ni decay can provide the strong emission in these band passes. Among different types of core-collapse SNe, SNe IIb/Ib/Ic (an explosion of a He or C+O star) are most promising. We estimate that the peak date in the high energy emission will be similar to (or a bit delayed than) that of SNe Ia (i.e., 2 - 3 weeks since the explosion). The smaller amount of $^{56}$Ni here (typically $\\sim 0.1 M_{\\odot}$), taking into account the delayed peak date, results in the expected horizon as $5 - 8$ Mpc (conservative and optimistic). Such nearby objects are rare, but expected to be discovered at a rate of a few in a decade according to the past statistics.    \\thanks KM thanks Franco Giovannelli and the organizers of Frascati Workshop 2013 for creating the friendly and stimulating atmosphere. The work by KM has been supported by WPI initiative, MEXT, Japan. The authors acknowledge financial support by Grant-in-Aid for Scientific Research from MEXT (22684012, 23340055, 23740141)."
    },
    "1402/1402.3874_arXiv.txt": {
        "abstract": "We present the results of our analysis of the RR Lyrae (RRL) variable stars detected in two transition-type dwarf galaxies (dTrans), ESO294-G010 and ESO410-G005 in the Sculptor group, which is known to be one of the closest neighboring galaxy groups to our Local Group. Using deep archival images from the Advanced Camera for Surveys (ACS) onboard the Hubble Space Telescope (HST), we have identified a sample of RR Lyrae candidates in both dTrans galaxies [219 RRab (RR0) and 13 RRc (RR1) variables in ESO294-G010; 225 RRab and 44 RRc stars in ESO410-G005]. The metallicities of the individual RRab stars are calculated via the period-amplitude-[Fe/H] relation derived by Alcock et al. This yields mean metallicities of $\\langle [Fe/H] \\rangle_{ESO294} = -1.77 \\pm 0.03$ and $\\langle [Fe/H] \\rangle_{ESO410} = -1.64 \\pm 0.03$. The RRL metallicity distribution functions (MDFs) are investigated further via simple chemical evolution models; these reveal the relics of the early chemical enrichment processes for these two dTrans galaxies. In the case of both galaxies, the shapes of the RRL MDFs are well-described by pre-enrichment models. This suggests two possible channels for the early chemical evolution for these Sculptor group dTrans galaxies: 1) The ancient stellar populations of our target dwarf galaxies might have formed from the star forming gas which was already enriched through ``prompt initial enrichment'' or an ``initial nucleosynthetic spike'' from the very first massive stars, or 2) this pre-enrichment state might have been achieved by the end products from more evolved systems of their nearest neighbor, NGC 55. We also study the environmental effects of the formation and evolution of our target dTrans galaxies by comparing their properties with those of 79 volume limited ($D_{\\odot} < $2 Mpc) dwarf galaxy samples in terms of the luminosity-metallicity relation and the {{\\sc H~i}\\/} gas content. The presence of these RRL stars strongly supports the idea that although the Sculptor Group galaxies have a considerably different environment from the Local Group (e.g. no giant host galaxies, loosely bound and very low local density), they share a common epoch of early star formation with the dwarf satellite galaxies in the Local Group. ",
        "introduction": "The chemical evolution of galaxies during their earliest stages provides important insights for the initial physical and environmental conditions of the galaxy formation process. The star formation history (SFH) of galaxies is one of the important factors determining the path of their chemical evolution because the number of stars formed (i.e. the star formation rate) with a given initial mass function (IMF) controls the rate of chemical enrichment (Matteucci 2012). One key observable signature of the SFH of a galaxy is the metallicity distribution function (MDF) of its stars. Thus, we can test possible scenarios for the early epochs of galaxy formation by comparing the observed MDF of the oldest stellar population in the galaxy with several analytical models of chemical evolution.  In this context, the Sculptor group dwarf galaxies provide an excellent test bed for investigating the early chemical evolution of dwarf satellite galaxies formed in low density environments (i.e. outskirts of dense galaxy clusters,  Local Group-like group environment, or field-like environment; see Section 4.4 of the present study). The Sculptor group, also known as the Sculptor filament, is one of the closest groups of galaxies to our Local Group (i.e. distance to the group center, $d_{gc} \\sim $3.9 Mpc; Karachentsev et al. 2003). Five bright late-type galaxies (NGC 55, 247, 253, 300, and 7793) and at least 16 dwarf satellites form a very loosely bound, field-like system stretched along the line of sight over $\\sim$5 Mpc. Our target galaxies, ESO294-G010 and ESO410-G005, are Phoenix-like transition-type dwarf galaxies (dTrans) that morphologically resemble the gas-poor dwarf spheroidal galaxies (dSphs) but contain detectable amounts of neutral hydrogen (H I) gas (i.e. $\\sim10^5 M_{\\odot}$; Bouchard et al. 2005; Jerjen, H., Freeman, K. C., \\& Binggeli, B. 1998), and show signs of recent star formation activity. Most recently, from their detailed analysis of the color-magnitude diagrams obtained using deep HST/ACS imaging for the five Sculptor group dwarf galaxies, Lianou et al. (2013; hereafter L13) found population gradients in ESO294-G010 and ESO410-G005 in the sense that young blue main sequence stars (age $<$ 100 Myr) and intermediate age asymptotic giant branch stars (age $\\sim$ 1-2 Gyr) are more centrally concentrated, while old horizontal branch stars (age $>$ 10 Gyr) tend to be more spatially extended.  The high quality of those HST/ACS images also allowed the first detection of a significant population of RR Lyrae variable star candidates beyond the Local Group. Da Costa et al. (2010) were able to discover numerous RR Lyrae candidates in both dTrans galaxies using a template light curve technique (Layden 1998). The presence of the RR Lyrae stars in ESO294-G010 and ESO410-G005 directly confirms the existence of ancient stellar populations with ages $>$ 10 Gyr and indicates that the Sculptor group dwarf galaxies share a common epoch of the earliest star formation with our Local Group galaxies even though these two groups of galaxies have quite different local density environments.  In addition to being a direct probe of ancient stellar populations, RR Lyrae stars provide a variety of utilities to investigate a number of important astrophysical applications; these stars are well-known reliable population II distance indicators. There is a correlation between the metallicities and the pulsation properties (i.e. periods and amplitudes) of the fundamental mode RR Lyrae stars (RR0 or RRab). Using this relation (Alcock et al. 2000; Sarajedini et al. 2006), we can calculate the metallicities of individual RRab stars, and thus construct the MDF for a purely old stellar population. Lastly, the intrinsic colors of RRab stars at their faintest luminosity are largely independent of their other physical properties [$(V-I)_{0,min}=$0.58$\\pm$0.02; Guldenschuh et al. 2005], therefore using their observed colors at minimum light we can also estimate the line-of-sight reddening.  The main goal of the present study is to examine the MDFs of legitimate RR Lyrae stars in the two Sculptor group dTrans galaxies, ESO294-G010 and ESO410-G005, in order to investigate the early chemical evolution of these systems, especially in light of the low density environments in which they reside. This paper is organized as follows. Section 2 provides a description of the data set and photometry process. In section 3 we report the results of our analysis including the general trends in the color-magnitude diagrams (CMDs) of the galaxies (Sec 3.1), the details of our RR Lyrae detection method and the pulsation properties of the RR Lyrae candidates (sec 3.2), the calculation of the metallicities of individual RR Lyraes (sec 3.3), and the distance measurements (sec 3.4). In section 4 based on our estimates in the previous sections, we present our in depth discussion on the star formation histories, early chemical evolutions, the luminosity-metallicity relations of the our target dTrans, and the environmental effects on the evolution of the near field dwarf galaxies. Finally in section 5 we present a summary of our results.  \\vskip 1cm ",
        "conclusions": "We have presented an extensive analysis of the RR Lyrae stars in two transition type Sculptor group dwarf galaxies, ESO294-G010 and ESO410-G005. Based on the properties of the RR Lyrae stars we present the following results.  \\begin{enumerate} \\item We have detected numerous RR Lyrae candidates in both ESO294-G010 ($N_{ab}=$219) and ESO410-G005 ($N_{ab}=$225). Based on the Bailey diagrams, the characteristics of the RR Lyrae stars in these two Sculptor group dTrans appear to follow the typical trend of the RR Lyrae populations of Local Group dwarf satellite galaxies.  \\item We construct the MDFs of the RRab stars using the period-amplitude-[Fe/H] relationship of Alcock et al. (2000). The mean metallicities estimated from the best-fit Gaussians to the MDFs are $<[Fe/H]>=$--1.77 $\\pm$ 0.03 (sem) for ESO294-G010, and $<[Fe/H]>=$-1.64 $\\pm$ 0.03 (sem) for ESO410-G005 respectively. Our results represent the metallicity values for purely ancient stellar populations in both dTrans galaxies.  \\item The distance of each dTran was calculated using the absolute V magnitudes of the RRab stars. We find $(m-M)_{0}=$ 26.40 $\\pm$ 0.07 mag for ESO294-G010 and $(m-M)_{0} =$ 26.33 $\\pm$ 0.07 mag for ESO410-G005. We have also calculated the distances of the two dTrans using the I-band magnitudes of TRGBs in order to check the validity of our RR Lyrae distance estimates [$(m-M)_{0,ESO294}=$26.37; $(m-M)_{0,ESO410}=$26.38]. Our distance estimates for ESO294-G010 using two independent methods (RRLs and TRGB) agree very well with each other. In the case of ESO410-G005, there is a $\\sim$ 0.1 mag difference but this is consistent  within the margin allowed by the errors.  \\item We have compared the RR Lyrae MDFs for ESO294-G010 and ESO410-G005 with several chemical evolution models. For both galaxies, the shapes of the RR Lyrae MDFs are nicely described by pre-enrichment models. This suggests two possible channels for the early chemical evolution of these Sculptor group systems: 1) The ancient stellar populations of our target dwarf galaxies might have formed from the star forming gas which was already enriched through ``prompt initial enrichment'' or an ``initial nucleosynthesis spike'' caused by the very first massive stars, or 2) This pre-enrichment state might have been achieved by the end products from more evolved systems of their nearest neighbor, NGC 55. We also fit a simple accretion model to the RR Lyrae MDFs of each galaxy. The gas infall hypothesis seems to explain the metal-poor tail of the RR Lyrae MDFs indicating that we cannot completely rule out the role of gas accretion on the early chemical evolution of the dTrans.  \\item The L-M relation of our target dTrans ESO294-G010 and ESO410-G005 follow the typical trend of other Local Group dTrans galaxies, such as Phoenix, Tucana, Cetus, LGS 3, KKR 25, DDO 210, and Antlia. Within a luminosity range of 5 $< log (L_{K}/L_{\\odot}) <$ 7.5, the L-M relation of the sample dTrans appear to show a closer resemblance to that of the canonical dSphs than dIrrs. This suggests that the bulk of the dTran galaxies can be considered as ``$present-day$ $progenitors$'' of the dSph galaxies if the recent star formation of dTrans galaxies has been largely suppressed by rapid gas loss (G03).  \\item Our examination of the {{\\sc H~i}\\/} gas content as a function of the tidal index $\\Theta_{5}$ for 79 volume limited ($D_{\\odot} < $2 Mpc) dwarf galaxies reveals a clear morphological segregation, in the sense that gas-deficient dSphs tend to be located in dense areas while most gas-rich dIrrs (except LMC/SMC pairs) are isolated. Similar to the dIrr galaxies, most dTran galaxies are also found to be isolated but in general they have lower {{\\sc H~i}\\/} gas fractions as compared to the dIrrs. Our analysis supports the idea that the morphology-density relation appears to be a ubiquitous phenomena in cluster environments regardless of their shapes and richness (van der Wel et al. 2010, and references therein), and the environmental effects on galaxy evolution operate in a similar fashion even in such low density environments of the Local Group or its neighboring groups. \\end{enumerate}"
    },
    "1402/1402.6726_arXiv.txt": {
        "abstract": "Pulsating stars in eclipsing binary systems play an important role in asteroseismology. The combination of their spectroscopic and photometric orbital solutions can be used to determine, or at least to constrain, the masses and radii of components. To successfully perform any seismic modelling of a star, one has to identify at least some of the detected modes, which requires precise time-series photometric and spectroscopic observations. This work presents a progress report on the analysis of two $\\beta$ Cephei-type stars in eclipsing binaries: HD\\,101794 (V916 Cen) and HD\\,167003 (V4386 Sgr). ",
        "introduction": "HD\\,101794 and HD\\,167003 have been observed at the South African Astronomical Observatory (SAAO) between May 2 and 19, 2009. The $UBVRI$ time-series photometry has been acquired using the UCT CCD detector at the 1.0-m telescope. Spectroscopic observations were carried out using the GIRAFFE {\\'e}chelle spectrograph at the 1.9-m telescope. We obtained spectra in the wavelength range between 4200 {\\AA} and 6900 {\\AA} with a resolution $R \\approx$ 39000. Stellar magnitudes were calculated using the program DAOPHOT II (\\cite[Stetson 1987]{stet87}). Spectra were wavelength-calibrated and extracted using the IRAF software package (\\cite[Tody 1993]{tody93}). The {\\'e}chelle orders were merged and normalised with a program of our own, and the radial velocities were calculated by cross-correlating the observed spectra with non-LTE models of \\cite[Lanz \\&~Hubeny (2007)]{lahu07} using the method of \\cite[Tonry \\&~Davis (1979)]{toda79}. ",
        "conclusions": "The results for HD\\,167003 are very encouraging. First of all, we have confirmed it is indeed a pulsating star in a binary system. In addition, we have confirmed that the orbital period amounts to 10.79824 d and is not twice as long, as initially suspected by \\cite[Pigulski \\&~Pojma{\\'n}ski (2008)]{pipo08}. This star seems to be a single-lined spectroscopic binary, and our modelling suggests that the orbit is close to circular.  \\cite[Pigulski \\&~Pojma{\\'n}ski (2008)]{pipo07} detected only the primary eclipse in the ASAS-3 photometry of HD\\,167003. Our initial hypothesis was that the lack of the secondary eclipse is caused by a highly eccentric orbit. In light of the results of our modelling, we now know this cannot be the case. This suggests that the secondary eclipse is very shallow, and that the contribution of the secondary component to the total flux is small. While our present photometry is more accurate than the ASAS-3 photometry used by the previous investigators, we are also unable to find the secondary minimum in our observations. This could be because the orbital period of HD\\,167003 is quite long, and our phase coverage is incomplete. However, we managed to detect hints of a minute reflection effect.  Once the orbital inclination is known from the light curve modelling, we will be able to use the mass function to constrain the masses of components. From the fact that at least two of the photometrically detected modes are also seen in the radial velocity data, it seems probable that the attempts to identify mode degrees from amplitude ratios and phase differences will be successful. Therefore, this star seems to be a very good candidate for asteroseismic analysis."
    },
    "1402/1402.0793_arXiv.txt": {
        "abstract": "We have evaluated the electron capture rates on $^{20}$Ne, $^{20}$F, $^{24}$Mg, $^{24}$Na and the $\\beta$ decay rates for $^{20}$F and $^{24}$Na at temperature and density conditions relevant for the late-evolution stages of stars with $M=8$--12~M$_\\odot$. The rates are based on recent experimental data and large-scale shell model calculations.  We show that the electron capture rates on $^{20}$Ne, $^{24}$Mg and the $^{20}$F, $^{24}$Na $\\beta$-decay rates are based on data in this astrophysical range, except for the capture rate on $^{20}$Ne, which we predict to have a dominating contribution from the second-forbidden transition between the $^{20}$Ne and $^{20}$F ground states in the density range $\\log \\rho Y_e (\\mathrm{g~cm}^{-3}) = 9.3$--9.6. The dominance of a few individual transitions allows us to present the various rates by analytical expressions at the relevant astrophysical conditions. We also derive the screening corrections to the rates. ",
        "introduction": "Electron captures on nuclei play a crucial role in the high-density environment of late-stage stellar evolution~\\cite{Bethe:1990,Langanke.Martinez-Pinedo:2003} with three important consequences. It reduces the pressure which the degenerate relativistic electron gas can supply against the gravitational contraction of the stellar core. Furthermore it cools the core environment as the neutrinos produced in the capture process can leave the star virtually unhindered (as long as the density is less than about $10^{11}$~g~cm$^{-3}$ and carry away energy. Finally electron captures change protons in the nucleus into neutrons and hence drive the stellar composition more neutron rich.  Improving on the pioneering work by Fuller, Fowler and Newman (FFN)~\\cite{Fuller.Fowler.Newman:1980,*Fuller.Fowler.Newman:1982a,*Fuller.Fowler.Newman:1982b,Fuller.Fowler.Newman:1985} and making use of advances in nuclear modeling and in computational hard- and software development, electron capture rates have been determined for sd-shell nuclei ($A=17$--39)~\\cite{Oda.Hino.ea:1994} and for pf-shell nuclei ($A=45$--64) \\cite{Caurier.Langanke.ea:1999,% Langanke.Martinez-Pinedo:2000,Langanke.Martinez-Pinedo:2001} based on large-scale shell-model diagonalization calculations. The reliability of the calculations benefitted also strongly from experimental data for the Gamow-Teller (GT$_+$) distribution in nuclei (e.g.~\\cite{frekers:2006,Fujita.Rubio.Gelletly:2011,Sasano.Perdikakis.ea:2011}) which determine the electron capture rates at the stellar conditions for which nuclei in the mass range $A=17$--64 dominate the stellar matter composition. Indeed a detailed comparison of stellar capture rates derived from experimental GT$_+$ distributions for all $pf$-shell nuclei, for which data exist, with modern shell-model rates convincingly validated the use of the latter in late-stage stellar evolution studies~\\cite{Cole.Anderson.ea:2012} (for applications and consequences see~\\cite{Janka.Langanke.ea:2007}). We note that diagonalization shell-model calculations are yet not globally feasible for the very neutron-rich nuclei with $A>64$ which dominate the electron capture at densities in excess of a few $10^{10}$~g~cm$^{-3}$~\\cite{Juodagalvis.Langanke.ea:2010} and the respective rates must be determined based on other approaches such as the Random Phase Approximation (RPA) with occupation numbers from Shell Monte Carlo~\\cite{Langanke.Martinez-Pinedo.ea:2003}, thermofield dynamics approach~\\cite{Dzhioev.Vdovin.ea:2010}, and finite-temperature Quasiparticle RPA~\\cite{Paar.Colo.ea:2009,Niu.Paar.ea:2011}.  While stellar electron captures usually occur on an ensemble of nuclei present in the matter composition, capture on the specific nuclei $^{20}$Ne and $^{24}$Mg has been identified as crucial for the core collapse of 8--12~M$_\\odot$ stars \\cite{Nomoto:1984,*Nomoto:1987,hillebrandt.nomoto.wolff:1984}. Stars in this mass range develop degenerate ONe or ONeMg cores which are driven towards collapse in a process dubbed electron capture supernova triggered by the loss of electron pressure support due to electron captures, mainly on the very abundant nuclear species $^{20}$Ne and $^{24}$Mg \\cite{Nomoto:1984,*Nomoto:1987,hillebrandt.nomoto.wolff:1984,Huedepohl.Mueller.ea:2010}. We note that 8--12~M$_\\odot$ stars crucially contribute to the nucleosynthesis of specific nuclides. Its role for the synthesis of r-process elements is currently controversially discussed~\\cite{Ning.Qian.Meyer:2007,Janka.Mueller.ea:2008}.  Simulations of late-stage evolution of 8--12~M$_\\odot$ stars and electron-capture supernovae usually adopt the weak-interaction rates, including those for electron capture on $^{20}$Ne and $^{24}$Mg, from the work of Oda et al. \\cite{Oda.Hino.ea:1994}. These authors made available rate tabulations for an extensive set of nuclei in the mass range $A=17-39$, however, on a rather sparse temperature-density grid which is argued to be insufficient for detailed studies of the evolution stage for which the weak rates are essential~\\cite{Jones.Hirschi.ea:2013}. Shell-model rates for electron captures on $^{20}$Ne and $^{24}$Mg had previously to the work by Oda \\emph{et al.}~\\cite{Oda.Hino.ea:1994} been calculated by Takahara \\emph{et al.}~\\cite{Takahara.Hino.ea:1989}. Importantly the calculations of Takahara \\emph{et al.} and of Oda \\emph{et al.} had been performed before the Gamow-Teller strength distributions for $^{20}$Ne and $^{24}$Mg have been determined by charge-exchange experiments. Due to the isospin symmetry of the two nuclei, this goal could not only be achieved by techniques which determine the GT$_+$ strength distribution using $(d,{}^2\\textrm{He})$ and $(t,{}^3\\mathrm{He})$ reactions, but also by those measuring the GT$_-$ distribution by $(p,n)$ and $(^3\\textrm{He},t)$ reactions.  (In the latter a neutron is changed into a proton supplying the information required for $\\beta^-$ decays). The availability of these data calls for a reevaluation of the electron capture rates which we will present in this manuscript.  Besides the incorporation of recent experimental GT$_+$ data, we improve the previous rates also in two other important aspects. At first we point to the relevance of the ground-state-to-ground-state transition in the capture on $^{20}$Ne which, although it is of forbidden nature, is likely to dominate the capture rate in the astrophysically relevant temperature-density range for the core evolution of 8-12 M$_\\odot$ stars. Secondly, we correct the capture rates for screening effects in the dense environment (which decrease electron capture rates, but increase the competing $\\beta$-decays). Our study is completed by a reevaluation of the rates for electron captures and $\\beta$ decays of $^{20}$F and $^{24}$Na, which are the daughters of the electron capture processes on $^{20}$Ne and $^{24}$Mg, respectively.  We note that the electron capture rate on $^{20}$Ne as well as the $^{20}$F $\\beta$-decay rate is dominated by a few transitions which are experimentally determined, except for the forbidden ground-state-to-ground-state transition for which only an upper limit is known. The dominance of  a few transitions allows us to present the rates by an analytical expression for the relevant temperature-density region which removes uncertainties associated with extrapolations required for rate tabulations provided on a grid. ",
        "conclusions": "We have calculated the rates for electron captures on $^{20}$Ne, $^{20}$F, $^{24}$Mg, $^{24}$Na and $\\beta$ decays of $^{20}$F and $^{24}$Na which are key quantities for studies of the late-time evolution of 8--12~M$_\\odot$ stars.  So far such late-time studies are based on the rate evaluations of~\\cite{Takahara.Hino.ea:1989} and~\\cite{Oda.Hino.ea:1994}. We have improved these rates in three important aspects. First, we have incorportated experimental data from either $\\beta$ decay or charge-exchange experiments which have not been available at the time when Takahara \\emph{et al.}~\\cite{Takahara.Hino.ea:1989} and Oda \\emph{et al.}~\\cite{Oda.Hino.ea:1994} did their work. In our study the recent experimental data are supplemented by Gamow-Teller transitions derived from large-scale shell model calculations, similar to the procedure in~\\cite{Takahara.Hino.ea:1989,Oda.Hino.ea:1994}. Importantly we find that nuclear physics input into the astrophysically relevant rates for electron captures on $^{20}$Ne and $^{24}$Mg and the competing beta decays of the respective daughters is completely based on experimental data. The exception is the electron capture on $^{20}$Ne in the density regime $\\log \\rho Y_e = 9.3$--9.7. As our second improvement we point out that at temperatures $T < 0.7 \\times 10^9$ the capture rate is likely to be dominated by the second-forbidden transition from the $^{20}$Ne ground state to the $^{20}$F. Experimentally only an upper limit exists, which, however, is of the order of typical second-forbidden transition strengths. While we have used the upper limit as an estimate for this transition in our present work a calculation of the transition with an appropriate method like the shell model or an experimental determination is highly desirable.  As the third improvement, we have corrected the various rates for medium-induced effects. Here we followed the formalism discussed in ref.~\\cite{Juodagalvis.Langanke.ea:2010} for electron captures and extended it to the treatment for $\\beta$-decays. The environment reduces the electron chemical potential and enhances (reduces) the reaction $Q$-value for electron captures ($\\beta$ decays). As a consequence electron capture rates are lower in dense astrophysical environments than for bare nuclei, while $\\beta$ decay rates are larger.  For the astrophysical conditions here, the medium corrections change the rates typically by a factor of order two. The effect is, of course, significantly larger at such densities where the rates change from dominance of a certain transition to another (as is the case in the weak processes here) as these transitions are extremely sensitive to the effective $Q$ values.  We note that medium effects should also have a significant effect on the densities at which so-called URCA pairs operate and influence the late-stage evolution of the stars. As $\\beta$ decay rates are enhanced and the competing electron capture rates are lowered, the medium modifiactions will move the operation of the URCA pairs to somewhat higher densities. As the shifts of the electron chemical potential and of the $Q$ values are of order 100~keV under the relevant density (and temperature) conditions encountered, we expect that the URCA pairs operate at densities which are about $0.1 \\times 10^9$~g~cm$^{-3}$ larger than found in calculations which do not consider medium corrections on the rate. Stellar evolution studies which investigate the impact of screening on the URCA pairs are needed.  We have presented analytical expressions for both electron capture and beta-decay rates that allow for an accurate description of these processes for conditions at which URCA pairs operate in both intermediate mass stars~\\cite{Tsuruta.Cameron:1970} and neutron stars~\\cite{Schatz.Gupta.ea:2014}.  Rate tables on fine grids in temperature and density in the ranges $\\rho Y_e = 10^8 - 10^{10}$~g~cm$^{-3}$ and $T=10^8 - 10^{10}$~K can be obtained by request from the authors."
    },
    "1402/1402.4045_arXiv.txt": {
        "abstract": " ",
        "introduction": "Scalar field models have played a vital role in cosmological theoretical studies in nearly half a century. Those assumed scalar fields appeared in different cosmological research aspects  to settle different cosmological problems \\cite{1}, such as to drive inflation, to explain a time variable cosmological $'$constant$'$ and so on. The scalar fields have played another important essential role in the past nearly fifteen years as a candidate of dark energy after the discovery of the accelerating expansion of universe. There are so many phenomenological dark energy models of scalar fields, such as quintessence, phantom, quintom and the scalar fields with non-canonical kinetic energy term (for a review, see\\cite{2,3}).  \\par To study the dynamical evolution of those scalar fields models and their cosmological implications with a phase-plane analysis is a very useful and common method(see Ref\\cite{add2} and recent papers, e.g. \\cite{Saridakis1, Saridakis2, Saridakis3, Saridakis4}). However, most of those works only focus on the quintessence models(including phantom quintessence and quintom) with unique exponential potential and tachyon models(including phantom tachyon) with inverse square potential, and correspondingly, the dynamical systems are two dimensional autonomous system with those special form of potentials(see the references cited in \\cite{6, 12}). Using a method which considers the potential related variable $\\Gamma$ as a function of another potential related variable $\\lambda$ (see Eq.(\\ref{eqsadd4}) for the definition of $\\Gamma$ and $\\lambda$)\\cite{26, 12, 6}, we are able to analyze the phase-plane of the dynamical systems of the quintessence and tachyon models with many different potentials. When the potentials are beyond the special type such as exponential or inverse square potentials, the dynamical systems consequently become a three-dimensional autonomous systems. This method is quite effective and powerful, it therefore has been generalized to several other cosmological contexts\\cite{cite61, cite62, cite63, cite65, cite66, cite67, cite68, cite69, 25, cite611}. However, there is very few work focusing on the dynamical behavior of the scalar field with a general modified kinetic term, such as k-essence ($L=V(\\phi)F(X)$) and general non-canonical scalar field ($L=F(X)-V(\\phi)$). Recently, Josue De-Santiago et. al analyzed the dynamical system of general non-canonical scalar field with the lagrangian $L=F(X)-V(\\phi)$ and studied the phase plane after a suitable choice of variables\\cite{20}. They obtained the three-dimensional autonomous system of this non-canonical scalar field after specifying the kinetic term as $F(X)=AX^{\\eta}$ and choosing the potential as $V(\\phi)=V_0(\\phi-\\phi_0)^{1/(1-\\Gamma)}$(i.e., the special case that potential related parameter $\\Gamma$ is a constant) and studied the critical points as well as their stability. \\par Motivated by the work, in this paper we try to extend our works in\\cite{6, 12} to give the three-dimensional autonomous dynamical systems for most of the popular scalar field dark energy models including (phantom) quintessence, (phantom) tachyon, k-essence and general non-canonical scalar field models. We will show that the three-dimensional autonomous systems of general non-canonical scalar field and k-essence will reduce to the quintessence and tachyon scalar field respectively. Not like the previous works, here we express the three dimensional autonomous systems from the trivial variables $(x, y, \\lambda)$ to the observable related variables $(w_{\\phi}, \\Omega_{\\phi}, \\lambda)$. It will be very convenient to investigate the dynamical properties of the autonomous system based on the observable related variables $w_{\\phi}$ and $\\Omega_{\\phi}$(see \\cite{7, 8, 9, 10, 12, 13} and a recent paper about the general property of dynamical quintessence field\\cite{add1}). Since the definition of the trivial variables $x$ and $y$ could vary with different scalar field models, while the observable related variables such as the equation of state $w_{\\phi}$ and the dark energy density parameter $\\Omega_{\\phi}$ are the same for different dark energy models. The paper is organized as follows. We firstly present the basic theoretical framework for (phantom) quintessence, (phantom) tachyon, k-essence and general non-canonical scalar field models in section 2, and try to give the relationships between those different scalar fields in this section. We then give the three dimensional autonomous dynamical systems for those scalar fields and switch the dynamical variables from $(x, y)$ to $(w_{\\phi}, \\Omega_{\\phi})$ in each subsection of section 3. Additionally, using the dynamical systems, we give the exact solution of $w_{\\phi}$ and $\\Omega_{\\phi}$ for a special case of tachyon model when the potential is chosen to be a constant in subsection 3.2. We show that the dynamical autonomous system of k-essence can reduce to tachyon model, and investigate another special case called kinetically driven quintessence with the lagrangian $p(X, \\phi)=f(\\phi)(-X+X^2)$ detailedly in subsection 3.3. In subsection 3.4, we show that the dynamical autonomous system of general non-canonical scalar field can reduce to quintessence and tachyon model respectively for some special cases. We also studied two special cases of purely kinetic united model $L = F(X)$ in detailed in this subsection. We try to give the cosmological implications of the three-dimensional dynamical autonomous system and present the conclusion in section 4. We proved in section 4 that the dark energy density parameter $\\Omega_{\\phi}$ would obey the same differential equation for all the non-coupled dark energy model under the GR frame. We also raise a question about the possibility of the chaotic behavior in the spatially flat single scalar field FRW cosmological models in the presence of ordinary matter. ",
        "conclusions": "\\par The main purpose of the paper is not to analyze the dynamical behavior about different scalar fields in detail, so we will not investigate the detailed critical points and their stable properties for each dynamical system. What we want to focus on is about the dynamical system itself. \\par First thing we want to emphasize is that, the dynamical variables ($x, y$) in the previous papers are about the scalar field and its first derivative. these dynamical variables have no direct cosmological meaning, and the autonomous system for $dx/dN$ and $dy/dN$ also varies from different models. However, it would be more convenient if we change the variables from $(x, y)$ to observable quantities $(\\gamma_{\\phi}, \\Omega_{\\phi})$. Firstly, $(\\gamma_{\\phi}, \\Omega_{\\phi})$ are directly related to the observable quantities and also about the the properties of dark energy. Analyzing the system based on $(\\gamma_{\\phi}, \\Omega_{\\phi})$, we can figure out how the equation of state of dark energy $w_{\\phi}$ and the density parameter $\\Omega_{\\phi}$ evolve. Secondly, though the form of autonomous system for $dx/dN$ and $dy/dN$ are completely different for different models, but we found that the function for $d\\Omega_{\\phi}/dN$ has the same expression in quintessence, tachyon, k-essence and general non-canonical scalar field model(i.e., Eq.(\\ref{eqs14}),Eq.(\\ref{eqs23}),Eq.(\\ref{eqs35}),Eq.(\\ref{eqs44})) as follows:  \\begin{equation}\\label{eqs47}\\frac{d \\Omega_{\\phi}}{dN}=3(\\gamma_b-\\gamma_{\\phi})\\Omega_{\\phi}(1-\\Omega_{\\phi})\\end{equation}  The only difference are the form of function $d\\gamma_{\\phi}/dN$. In fact, we can prove that Eq.(\\ref{eqs47}) holds for all non-coupled dark energy models as long as they satisfy the following equations: \\begin{equation}\\label{eqs48} H^2=\\frac{1}{3M^2_{pl}}(\\rho_b+\\rho_{de}), ~ \\dot{\\rho}_i+3H(p_i+\\rho_i)=0\\end{equation}  where subscript $i$ denotes each energy component such as dark energy, matter or radiation. If we set $\\gamma_{de}=w_{de}+1=p_{de}/\\rho_{de}+1$ and $\\Omega_{de}=\\rho_{de}/3M^2_{pl}H^2$, we can obtain following equation from Eqs.(\\ref{eqs48}):  \\begin{equation}\\label{eqs49}\\frac{d \\Omega_{de}}{dN}=3(\\gamma_b-\\gamma_{de})\\Omega_{de}(1-\\Omega_{de})\\end{equation}  Eq.(\\ref{eqs49}) is the same with Eq.(\\ref{eqs47}), and also the same with Eq.(\\ref{eqs14}), Eq.(\\ref{eqs23}), Eq.(\\ref{eqs35}) and Eq.(\\ref{eqs44}). So Eq.(\\ref{eqs49}) is a very common form for the evolution of $\\Omega_{de}$ of various different models. For example, the authors got the same equation even for the purely kinetic coupled gravity model which modified the standard general relativity action through the addition of a coupling between functions of the metric and kinetic terms of a free scalar field\\cite{24}(Eq.(28) in this paper and $\\gamma_b$ is taken as 1).  \\par Eq.(\\ref{eqs49}) holds for the different scalar field models mentioned in this paper and even all the non-coupled dark energy models which satisfy Eq.(\\ref{eqs48}), so we can conclude that there are only three possible destinies(three types of critical points) for $\\Omega_{de}$ to be in these models. The value of $\\Omega_{de}$ will be determined by the stable properties of these three types of critical points, and the stable properties depend on different models and the value of parameters in different models.  Two of them are the cases that $\\Omega_{de}=0$ or $\\Omega_{de}=1$ which are completely opposite destinies. These two solutions correspond to the universe completely dominated by the scalar field or by the barotropic fluid. The last destiny corresponds to the case of $\\gamma_{de}=\\gamma_b$, and the value of $\\Omega_{de}$ will vary from different models and be determined by other equations in the dynamical system. Generally we can  obtain $0\\leq\\Omega_{de}\\leq1$ and so this solution is called the scaling solution. However, for the scaling solution, the equation of state of dark energy $w_{de}$ is the same as the equation of state of barotropic fluid $w_b$, so there is no accelerating expansion. Since the observation suggested that we are living in a accelerated expanding  universe with $\\Omega_{de}\\sim0.7$, none of these three destinies correspond to the present universe we observed. If we try to construct a successful theoretical model which has a dynamical stable solution corresponding to present observational universe(namely a scaling attractor with $\\Omega_{de}\\sim 0.70$ and $\\gamma_{de}\\sim 0.1$ in agreement with observations) to solve or at least alleviate the cosmological coincidence problem without fine-tunings, we must consider the interaction between dark energy and other barotropic fluids(see \\cite{25} for such model). This result is valid for not only all the non- coupled dark energy models, but also for many modified gravity models as long as the energy density and the pressure of dark energy or effective dark energy satisfy the continuity equation Eq.(\\ref{eqs48}).  \\par Another important thing we want to emphasize is that, it is more reasonable and more scientific to investigate the dynamical behaviors of a dark energy model under the three-dimensional autonomous system rather than the two-dimensional system.  Firstly, the two-dimensional dynamical autonomous system is just a specific case when the potential takes a special form. if we want to completely study the general dynamical properties of a dark energy model, we need to study the system beyond a special potential. Then we can find more critical points than the ones found in a two-dimensional system. We therefore are able to analyze which critical points are possessed by a class of dark energy models and which ones exist only due to the concrete potentials. The method studying the three-dimensional dynamical autonomous system beyond one special potential is originated for the quintessence \\cite{6, 26} and then developed to other dark energy models \\cite{cite61, cite62, cite63, cite65, cite66, cite67, cite68, cite69, 25, cite611}. Here we extend this method to the more general scalar field models in Section 3. Secondly, from the viewpoint of chaos theory, the dynamical properties of three-dimensional autonomous system is more fruitful than the two-dimensional system. According to the Poincar$\\acute{e}$ -Bendixson theorem, chaos does not exist in any two-dimensional autonomous dynamical system\\cite{27, 28} but could be possible in three-dimensional autonomous dynamical system. For a number of three-dimensional systems, such as the famous three-dimensional Lorenz equations which is a model describing the atmospheric convection \\cite{29}, there exist chaos for certain values of the parameters. \\par The studies of chaotic dynamics in cosmological models has a long story. Chaotic properties had reported in spatially closed scalar field FRW cosmological models \\cite{30, 31, 32, 33, 34, 35}, spatially flat FRW cosmological model with two or more scalar fields \\cite{36,37}, Bianchi IX universe \\cite{38, 39}, Bianchi I universe\\cite{40} and the mixmaster universe\\cite{41}.  It would be very interesting and also a big challenge for the theoretical study of dark energy if the dynamical systems we consider here(i.e., Eqs.(\\ref{eqs14}-\\ref{eqs16}),Eqs.(\\ref{eqs23}-\\ref{eqs25}),Eqs.(\\ref{eqs35}-\\ref{eqs37}) and Eqs.(\\ref{eqs44}-\\ref{eqs46}) ) exist the chaotic properties. Then the evolution of $\\Omega_{\\phi}$ and $\\gamma_{\\phi}$ will be very sensitive to the initial condition, and therefore predicting their evolution in future becomes totally impossible. However, it is proved that there is no chaotic behavior in spatially flat single scalar field FRW cosmological models\\cite{42, 43}. Since for the spatially flat case with $k=0$, the dynamical system can be described by a three-dimensional autonomous system with a set of variables $(H, \\phi, \\dot\\phi)$ under a Hamiltonian constraint, so the dynamical system is actually a two-dimensional autonomous system($a$ and $\\dot a$ appear only in the combination $H=\\dot a/ a$)\\cite{44}. We know that for the two-dimensional autonomous systems, there are no enough degrees of freedom to exist chaos, so this proved no-chaotic dynamics in the spatially flat scalar field FRW cosmological model. However, we noted that this result is obtained in the absence of matter and radiation. This may be the case in the very early time when our universe is undergoing an inflation era and completely dominated by the scalar field. However, for the study of dark energy of late-time cosmic acceleration, the component of matter is comparable with the density of dark energy and should not be ignored when we investigate the dynamical behavior of scalar field. In the presence of matter, scale factor $a$ will reappear in the dynamical system beside the variables $(H, \\phi, \\dot\\phi)$, and then the system can not be reduced to two-dimensional autonomous dynamical system any more. So here we argue that it is still possible for the chaotic behavior in spatially flat single scalar field FRW cosmological models in the presence of matter. It is very like the case of spatially non-flat($k\\neq 0$) single scalar field FRW cosmological models where the dynamical system can not be reduced to two-dimensional autonomous dynamical system too. What we argued here is also supported by the equations in section 3( i.e., Eqs.(\\ref{eqs14}-\\ref{eqs16}), Eqs.(\\ref{eqs23}-\\ref{eqs25}), Eqs.(\\ref{eqs35}-\\ref{eqs37}) and Eqs.(\\ref{eqs44}-\\ref{eqs46})), which described three-dimensional autonomous dynamical systems. However, we are not sure whether there truly exist the chaotic behavior in spatially flat scalar field FRW cosmological models now, to find the chaotic behavior is beyond the scope of this paper, we will investigate it in future."
    },
    "1402/1402.3948_arXiv.txt": {
        "abstract": "The Galactic microquasar GRS 1915+105 exhibits various types of light curves. There is, however, no understanding of when a certain type of light curve will be exhibited and only in a handful of cases, the transitions from one type to another have actually been observed. We study the detailed spectral properties in these cases to show that that different classes have different ratio of the power-law photon and the blackbody photon. Since the power-law photons are from the Compton cloud, and the intensity of the power-law photon component depends on the degree of interception of the soft photons by the Compton cloud, we conclude that not only the accretion rate, but the accretion flow geometry must also change during a class transition. ",
        "introduction": "Indian X-ray Astronomy Experiments (IXAE) payload on board IRS-P3 satellite data of the enigmatic stellar mass black hole binary GRS 1915+105 (Harlaftis \\& Greiner, 2004) clearly indicated that GRS 1915+105 can have various types of light curves (Yadav et al., 1999; Rao, Yadav \\& Paul, 2000; Chakrabarti \\& Nandi, 2000 and references therein; Naik et al. 2002a). So far, a total of thirteen types of light curves have been detected. However, there are only a handful of cases when an actual transition from one type to another has been reported (Chakrabarti et al., 2004; Chakrabarti et al. 2005, hereafter Paper I and Paper II respectively), and those too using the IXAE data which had only two energy channels. It was possible to see the transitions because the IXAE payload observed the same object, namely, GRS 1915+105, for several orbits. Specifically, Papers I and II showed that in a matter of two to three hours indicating that it is not the viscous time scale of the Keplerian disk, but the free-fall time scale of the low-angular momentum halo which decides the variability class transition.  In the present paper, we study the spectra of the classes which showed the transitions and found that one of the parameters which could be responsible for the class transition is the geometry of the Compton cloud. There are various models of this cloud, including a Keplerian  disk with corona (Haardt \\& Maraschi, 1993) and the unstable inner edge of the standard disk (Kobayashi et al. 2003). However, we shall use a physically reasonable solution of the two component advective flow (TCAF) model (Chakrabarti \\& Titarchuk, 1995, hereafter CT95), where, the Compton cloud is made up of the post-shock region of the sub-Keplerian component which includes the jet and the outflows as well. In the literature many works are present which discusses the quasi-periodic oscillations of black hole candidates (e.g., Strohmayer, 2001; Wagoner, Silbergleit,  Ortega-Rodr\u00edguez, 2001; Stella \\& Vietri, 1999; Abramowicz, Kluzniak, 2001; Stuchl\u00edk, Slan\u00fd, T\u00f6r\u00f6k, 2007; Kato, 2008; Blaes, Sramkova, Abramowicz, Kluzniak \\& Torkelson, 2007). However, we shall concentrate only on the spectral properties of GRS 1915+105 in the current paper.  The plan of this {\\it letter} is the following: we briefly present the data analysis procedure in our next section. In \\S 3, we study the nature of the components of the spectra, especially the black body and the power-law components. We compute the efficiency of Comptonization by taking the ratio of the photon numbers in those components. We show that indeed, the average ratio varies from class to class. Finally, in \\S 4, we discuss the physics behind such changes in relation to the accretion flow dynamics. ",
        "conclusions": "In this paper, we have analysed the classes of GRS 1915+105 which have demonstrated class transitions. Purely in a model independent way, we computed the ratios of the power-law photons and the soft photons coming from the Keplerian disk in all these classes. We ignore very high energy photons from HXETE since its contribution is less than a percent and excluding this does not change our conclusion. Since the power-law photons are are believed to be formed due to inverse Comptonization of the injected soft photons, the only way the ratio can change is to change the geometry of the Compton cloud. In CT95 model, the puffed-up part of the low angular momentum flow (CENBOL) remains big in harder (burst-off) states and collapses in softer  (burst-on) states. Thus we naturally see that the average Comptonizing efficiency in burst-off states is higher than that in burst-on states. Not only that, all the observed transitions are found to reside side-by-side in the class vs. CE plane. In other words, for a transition, the average CE changes smoothly. This implies that the geometry of the Compton cloud also varies from class to class. The geometry of a Compton cloud can depend on several physical conditions. In CT95, the CENBOL collapses when the accretion rate of the Keplerian disk is increased. This in softer state, CE is low. When the disk rate is very low, the CENBOL cannot be cooled down. It remains big, as in an ion-pressure supported torus (Rees et al. 1982), and thus the interception of soft photons could be large, thereby increasing CE. We believe that this interpretation can be used even for classes not discussed in this paper. This work is being done and the results would be presented shortly.  Recently,  Remillard and McClintock. (2006) presented evidances that  perhaps GRS 1915+105 is an extremely rotating Kerr black hole. However, results discussed in this paper are directly the analysis of the observational data and the interpretation of variation of the Comptonizing efficiency is totally model independent. In other words, our conclusion of the variation of the geometry from one class to another remains valid even when the spin of the black hole is extreme.  The work of P. S. Pal is supported by CSIR fellowship."
    },
    "1402/1402.1456_arXiv.txt": {
        "abstract": "Far-infrared and submillimeter wavelength surveys have now established the important role of dusty, star-forming galaxies (DSFGs) in the assembly of stellar mass and the evolution of massive galaxies in the Universe.  The brightest of these galaxies have infrared luminosities in excess of 10$^{13}$ L$_{\\odot}$ with implied star-formation rates of thousands of solar masses per year.  They represent the most intense starbursts in the Universe, yet many are completely optically obscured. Their easy detection at submm wavelengths is due to dust heated by ultraviolet radiation of newly forming stars.  When summed up, all of the dusty, star-forming galaxies in the Universe produce an infrared radiation field that has an equal energy density as the direct starlight emission from all galaxies visible at ultraviolet and optical wavelengths. The bulk of this infrared extragalactic background light emanates from galaxies as diverse as gas-rich disks to mergers of intense starbursting galaxies. Major advances in far-infrared instrumentation in recent years, both space-based and ground-based, has led to the detection of nearly a million DSFGs, yet our understanding of the underlying astrophysics that govern the start and end of the dusty starburst phase is still in nascent stage.  This review is aimed at summarizing the current status of DSFG studies, focusing especially on the detailed characterization of the best-understood subset (submillimeter galaxies, who were summarized in the last review of this field over a decade ago, \\citealt{blain02a}), but also the selection and characterization of more recently discovered DSFG populations.  We review DSFG population statistics, their physical properties including dust, gas and stellar contents, their environments, and current theoretical models related to the formation and evolution of these galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3723_arXiv.txt": {
        "abstract": "The Survey of \\HI\\ in Extremely Low-mass Dwarf galaxies (SHIELD) is an on-going multi-wavelength program to characterize the gas, star formation, and evolution in gas-rich, very low-mass galaxies. The galaxies were selected from the first $\\sim10$\\% of the \\HI\\ ALFALFA survey based on their inferred low \\HI\\ mass and low baryonic mass, and all systems have recent star formation. Thus, the SHIELD sample probes the faint end of the galaxy luminosity function for star-forming galaxies. Here, we measure the distances to the 12 SHIELD galaxies to be between $5-12$ Mpc by applying the tip of the red giant method to the resolved stellar populations imaged by the Hubble Space Telescope. Based on these distances, the \\HI\\ masses in the sample range from $4\\times10^6$ to $6\\times10^7$ \\msun, with a median \\HI\\ mass of $1\\times 10^7$ \\msun. The TRGB distances are up to 73\\% farther than flow-model estimates in the ALFALFA catalog. Because of the relatively large uncertainties of flow model distances, we are biased towards selecting galaxies from the ALFALFA catalog where the flow model underestimates the true distances. The measured distances allow for an assessment of the native environments around the sample members. Five of the galaxies are part of the NGC~672 and NGC~784 groups, which together constitute a single structure. One galaxy is part of a larger linear ensemble of 9 systems that stretches 1.6 Mpc from end to end. Two galaxies reside in regions with $1-4$ neighbors, and four galaxies are truly isolated with no known system identified within a radius of 1 Mpc. ",
        "introduction": "} Historically, optical surveys have cataloged the majority of dwarf galaxies in the nearby universe. However, optical identification of systems populating the extremely low-mass end of the galaxy luminosity function are hampered by the intrinsic faintness and small angular size of such systems. These low luminosity systems dominate the local population of galaxies and are important tracers of the distribution of mass. Outside of a dense group or cluster environment where low mass galaxies do not repeatedly experience ram pressure and/or tidal stripping of their ISM, populations of such extremely low mass dwarfs are expected to be predominantly gas-rich. Thus, large \\HI\\ surveys offer an opportunity to identify and catalogue such systems in a statistically complete volume.  The Arecibo Legacy Fast ALFA (ALFALFA) survey \\citep{Giovanelli2005, Haynes2011} is a blind extragalactic \\HI\\ survey to map the nearby \\HI\\ universe over 7000 sq. degrees of high Galactic latitude sky. With an \\HI\\ mass detection limit as low as $10^6$ \\msun\\ for galaxies in the local universe and $10^{9.5}$ \\msun\\ at the survey velocity limit of $z\\sim0.06$, the ALFALFA survey was designed to populate the faint end of the \\HI\\ mass function over a cosmologically significant volume. The low mass end of the \\HI\\ mass function has been sampled by previous studies including the wide area \\HI\\ Parkes All-Sky Survey \\citep[HIPASS;][]{Barnes2001, Meyer2004, Wong2006}. However, ALFALFA has higher angular and spectral resolution and higher sensitivity. Thus, the ALFALFA catalogue provides a representative sampling of the \\HI\\ mass function to lower masses in a much larger volume.  The Survey of \\HI\\ in Extremely Low-mass Dwarf systems \\citep[SHIELD;][]{Cannon2011} was designed to probe the early release of the ALFALFA dataset in a systematic investigation of nearby galaxies with HI masses $\\ltsimeq 10^7$ \\msun\\ outside the Local Group. From the first $\\sim10$\\% of the processed ALFALFA survey data, 12 systems were selected that had gas mass estimates between $10^6-10^7$ \\msun\\ based on \\HI\\ line widths and flow-model distances, and overall low masses based on an \\HI\\ full width at half maximum $< 65$ km s$^{-1}$ which discriminated against gas-poor massive galaxies. The high positional accuracy of the ALFALFA survey (i.e., better than 20$\\arcsec$) facilitated the identification of optical counterparts in SDSS imaging, revealing young, blue stellar populations in each galaxy. These low mass systems have retained \\HI\\ mass reservoirs of $10^6-10^7$ \\msun\\ over a Hubble time, apparently evolving in relative quiescence, and all have recent star formation activity, consistent with all previous observations of dwarf irregular galaxies \\citep[e.g.,][]{Hunter1982}.  Once the presence of stellar components was confirmed with SDSS imaging, follow-up observations became possible in order to obtain accurate distance measurements and to probe the evolutionary histories of the galaxies. Accurate distances to these very low mass galaxies are important for a number of reasons. First, accurate distances to galaxies reduce uncertainties in the \\HI\\ mass function made possible by the ALFALFA survey. Second, distance measurements enable distance dependent analyses of the individual systems to be placed on an absolute scale. For example, the distance measurement of 1.72$^{+0.14}_{-0.40}$ to a newly discovered, extremely metal-deficient galaxy from the ALFALFA survey, Leo~P, allows interpretation of the star formation and evolutionary processes to be placed on solid ground \\citep{Giovanelli2013, Rhode2013, Skillman2013, McQuinn2013}. Third, combining \\HI\\ velocity measurements and accurate distances allows an assessment of environment. Most dwarf galaxies are thought to exist in some form of association with other galaxies \\citep{Tully2006}, with truly isolated galaxies the exception even in low density environments. The Nearby Catalog of Galaxies \\citep{Karachentsev2013} compiles measurements of 869 galaxies within 11 Mpc, including calculations of tidal indices that parameterize the likelihood of neighboring galaxies being kinematically bound in groups. Clearly, such work depends critically on accurate distance measurements. Finally, \\HI\\ velocity measurements and accurate distances enable low mass galaxies to be used to help trace the distribution of matter in larger low-density galaxy volumes. Significant effort has been made to map the distribution of galaxies locally. The very local Hubble flow ($0.7$ Mpc $<$ D$_{MW} < 3.0$ Mpc) has been mapped with increasing accuracy, measuring the expansion of the universe outside the boundaries of the Local Group \\citep{Karachentsev2009}. The Extragalactic Distance Database \\citep{Tully2009}, the Cosmicflow program \\citep{Courtois2011a, Courtois2011b, Tully2012}, and the Cosmicflows$-$2 program \\citep{Tully2013} are working to expand this analysis out to larger distances. The population of low mass galaxies with accurate distances from the richly populated ALFALFA data set can make an important contribution to mapping the mass in the Local Volume (LV; defined by an approximate radius of 11 Mpc).  Here, we present the distance measurements to the 12 SHIELD galaxies based on Hubble Space Telescope (HST) optical imaging and the tip of the red giant branch (TRGB) standard candle method \\citep[e.g.,][]{Mould1986}. The paper is organized as follows. \\S2 describes the sample and summarizes the observations and data processing. \\S3 describes the distance determination methods and results. \\S4 maps the native environment around the SHIELD galaxies including all known galaxies within a radius of 1 Mpc of each system. \\S5 summarizes our conclusions. Future work on the SHIELD sample will include an investigation of their star formation histories (McQuinn et~al., in preparation), metallicity measurements based on nebular abundance analyses (Haurberg et~al., in preparation), and a study of the gas kinematics from VLA observations (Cannon et~al., in preparation). ",
        "conclusions": "} We have used optical imaging of resolved stellar populations obtained from the $HST$ to measure the TRGB distance to 12 low-mass, gas-rich galaxies in the SHIELD program. The distances to the galaxies range from $5 - 12$ Mpc. The CMDs show a population of young main sequence stars indicating these systems have experienced recent star formation. A subset of the sample also shows red and blue helium burning sequences suggesting the recent star formation has been on-going for $\\gtsimeq$ 100 Myr. The star formation histories will be investigated in detail in a future paper (McQuinn et al. in prep). Five galaxies are part of two galaxy groups, NGC~784 and NGC~672, which are likely a single, bound kinematic structure. AGC~174585 is located $\\sim1.0$ Mpc away from the previously identified association of galaxies, $14+19$, but is likely unbound to this structure. Similarly, AGC~731457 lies 0.9 Mpc from another system, DDO~83. AGC~749241 is situated in a more populated volume, forming part of 1.6 Mpc linear structure with six other galaxies. Similarly to the NGC~672 and NGC~784 groups, these galaxies are likely bound, lying along a dark matter filament. The nearest neighbor identified is $\\sim0.5$ Mpc away. Four galaxies, AGC~174605, AGC~182595, AGC~748778, and AGC~749237 appear to be truly isolated with no known neighbors identified within a radius of 1 Mpc, although three of these galaxies are at relatively large distances (i.e., $\\gtsimeq10$ Mpc) where our knowledge of the galaxy population is incomplete.  Based on the distances and the \\HI\\ measurements from the ALFALFA survey, the SHIELD galaxies have \\HI\\ masses ranging from $4\\times10^6$ to $6\\times10^7$ \\msun, with a median \\HI\\ mass of $1\\times 10^7$ \\msun\\ and 5 galaxies with an \\HI\\ mass below $10^7$ \\msun. Thus, the SHIELD program probes the low mass regime of gas-rich galaxies in a larger cosmological volume than previous surveys. As these 12 galaxies were selected from the first $\\sim10$\\% of the reduced ALFALFA database, it suggests that a large sample of low luminosity galaxies in this mass range exists in the nearby universe. Nominally, one might expect the total yield from the ALFALFA survey to be 10$\\times$ greater. However, investigation of the first 40\\% of the ALFALFA catalogue \\citep{Haynes2011} suggests the true yield of low luminosity galaxies with low \\HI\\ velocities and inferred masses is likely more than 12$\\times$ greater."
    },
    "1402/1402.4429_arXiv.txt": {
        "abstract": "{ Our understanding of radiative feedback and star formation in galaxies at high redshift is hindered by the rarity of similar systems at low redshift. However, the recently identified Green Pea (GP) galaxies are similar to high-redshift galaxies in their morphologies and star formation rates and are vital tools for probing the generation and transmission of ionizing photons. The GPs contain massive star clusters that emit copious amounts of high-energy radiation, as indicated by intense \\oiii~$\\lambda$5007 emission and \\heii~$\\lambda$4686 emission. We focus on six GP galaxies with high ratios of \\oiii~$\\lambda\\lambda$5007,4959/\\oii~$\\lambda$3727 $\\sim$10 or more. Such high ratios indicate gas with a high ionization parameter or a low optical depth. The GP line ratios and ages point to chemically homogeneous massive stars, Wolf-Rayet stars, or shock ionization as the most likely sources of the \\heii~emission. Models including shock ionization suggest that the GPs may have low optical depths, consistent with a scenario in which ionizing photons escape along passageways created by recent supernovae. The GPs and similar galaxies can shed new light on cosmic reionization by revealing how ionizing photons propagate from massive star clusters to the intergalactic medium. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5934_arXiv.txt": {
        "abstract": "\\noindent We study structure and dynamics of turbulent photospheric magnetic field in active region NOAA 11158 by characterizing spatial and temporal scaling properties of the line-of-sight (LOS) component. Using high-resolution high-cadence LOS magnetograms from SDO/HMI, we measured power-law exponents $\\alpha$ and $\\beta$ describing Fourier power spectra in wavenumber ($k$) and frequency ($f$) domains and investigated their evolution during the passage of the active region through the field of view of HMI. The flaring active region NOAA 11158 produces a one-dimensional spatial power spectral density that follows approximately a $k^{-2}$ power law -- a spectrum that suggests parallel MHD fluctuations in an anisotropic turbulent medium. In addition, we found that the values of $\\alpha$ capture systematic changes in the configuration of LOS photospheric magnetic field during flaring activity in the corona. Position-dependent values of the temporal scaling exponent $\\beta$ showed that, on average, the core of the active region scales with $\\beta >$ 3 surrounded by a diffusive region with an approximately $f^{-2}$-type spectrum. Our results indicate that only about 1 - 3 \\% of the studied LOS photospheric magnetic flux displays $\\beta\\approx\\alpha$, implying that Taylor's hypothesis of frozen-in-flow turbulence is typically invalid for this scalar field in the presence of turbulent photospheric flows. In consequence, both spatial and temporal variations of the plasma and magnetic field must be included in a complete description of the turbulent evolution of active regions. ",
        "introduction": "Solar active regions (ARs) are the central building blocks in the path to understanding the drivers of space weather. Major solar flares and coronal mass ejections (CME) originate from active regions, where strong ($\\approx 10 ^{3}$ gauss (G)) and complex magnetic field structures can accumulate sufficient free energy to power energetic eruptions. The high degree of complexity in terms of topology and spatial distribution present in the photospheric magnetic field appears to emerge from a turbulent photospheric plasma state (\\cite{2005ApJ...629.1141A}, \\cite{2010ApJ...720..717A}). In this state, field emergence, fragmentation, and dissipation associated with turbulent flows lead to highly irregular spatio-temporal distribution of the magnetic field. In a simple way, we can view an AR as a system that takes the magnetic field and evolves it into an unstable non-potential configuration by non-linear shear and stress. For this system to return to a lower-energy state, the excess free energy must be released in a bursty event in the corona while electrical currents are dissipated and potential field configuration is restored (\\cite{2011LRSP....8....6S}).  Non-linear dynamical processes, such as turbulence, are often studied using a description that involves statistical momenta of the turbulent field. For instance, in hydrodynamics (HD) and magnetohydrodynamics (MHD), kinetic and/or magnetic energy injection, transfer, and dissipation processes in turbulent flows are understood in terms of the scale-free behavior of their Fourier spectrum (\\cite{1993noma.book.....B}, \\cite{2011soca.book.....A}), which follows a power-law distribution in space and time. Kolmogorov's 5/3 law is a classic example of these phenomena \\cite{1941DoSSR..30..301K}. In the case of the photospheric magnetic field, statistical parametric analyses have been performed with the aim of quantifying the complexity present in the field (see \\cite{2005ApJ...629.1141A}; \\cite{2010AdSpR..45.1067M}). However, only recently, when better and more accurate measurements of the photospheric magnetic field have become available, a more coherent picture of its complexity has started to emerge (see \\cite{2005ApJ...629.1141A}).  Previous studies on the complexity in the photospheric magnetic field in ARs can be divided into two categories: (1) analysis of physical and statistical parameters of the magnetic field such as the effective connected magnetic field strength (\\cite{2007ApJ...661L.109G}, \\cite{2008GeoRL..3506S02G}), the strong gradient length (\\cite{2002ApJ...569.1016F}, \\citeyear{2003JGRA..108.1380F}), or the statistical momenta of the field spatial distribution (\\cite{2003ApJ...595.1277L}, \\citeyear{2003ApJ...595.1296L}; \\cite{2008ApJ...688L.107B}), and (2) description of magnetic structures based on transformations of the LOS component such as the spatial power scaling exponent (Fourier analysis; \\cite{2005ApJ...629.1141A}; \\cite{2010ApJ...720..717A}) or fractal dimension (wavelet analysis; \\cite{2010AdSpR..45.1067M}). For example, \\cite{2007ApJ...661L.109G} defined the AR effective connected magnetic field strength $B_{\\rm eff}$ as a measure of magnetic field complexity. The $B_{\\rm eff}$ parameter accounts for the connectivity of individual photospheric magnetic flux concentrations; therefore its value depends on the spatial distribution of the flux concentrations. Values of $B_{\\rm eff}$ were measured using LOS magnetograms averaged over 12 h -- a cadence too low in order to capture transient phenomena of magnetic concentrations, which encompass a wide range of temporal scales (\\cite{2012ApJ...748...60U,uritsky13}).  \\cite{2005ApJ...629.1141A} analyzed a sample of ARs using photospheric magnetic data from the {\\it Michelson Doppler Imager} (MDI; \\cite{1995SoPh..162..129S}) instrument onboard SoHO and the Digital Magnetograph (DMG) located at the Big Bear Solar Observatory. This study was focused on measuring power-law scaling exponents of the spatial power spectral densities. \\cite{2005ApJ...629.1141A} concluded that the derived exponents described the scale-free behavior of the magnetic field and served as indicators for differentiating between ARs that are prone to produce flaring activity and those that are flare-quiet.  A common approach in the studies mentioned above was to quantify the complexity present in the instantaneous spatial distribution of the photospheric field and then to observe subsequent time evolution of the spatial parameter. Consequently, spatial and temporal domain analyses have been conducted for the most part in an independent fashion. A question that then naturally arises concerns the coupling between these two domains: is there a way to link the spatial and temporal variations? The first step to study this coupling is to verify the possible validity of Taylor's hypothesis of frozen-in-flow turbulence (\\cite{Taylor18021938}) for the photospheric plasma. If the hypothesis is valid, determining the (temporal) spatial scaling ({\\it e.g.} \\cite{2002ApJ...577..487A}) is sufficient since the (spatial) temporal scaling is constrained to be identical. On the other hand, if the hypothesis is not valid, both spatial and temporal scaling must be considered in the analysis in order to provide a complete picture of the state of the photospheric magnetic field and plasma. In this report, we will demonstrate that the latter is indeed the case for NOAA AR 11158.  In this paper, we extend previous studies of solar AR magnetic field complexity by addressing both spatial and temporal variability of the LOS photospheric magnetic field across a wide range of scales. By constructing a more comprehensive picture of the turbulent spatio-temporal dynamics in the AR photospheric magnetic field, we will provide new information that can help to better understand, for example, magnetic energy release signatures in the photosphere and the coupling between the photosphere and corona. In Section 2, we describe the analyzed set of LOS magnetograms and the active region to which they belong, NOAA 11158. Section 3 explains the data analysis and discusses the results. We explain the method of measuring the power-law exponents in two separate subsections: the spatial scaling analysis (Section 3.1) and the temporal scaling analysis (Section 3.3), both based on the Fourier transform of LOS magnetic field. Our main results are reported and discussed in Sections 3.2 and 3.4, while in Section 4 we draw conclusions and outline future work. ",
        "conclusions": "We have investigated the turbulent state of the line-of-sight photospheric magnetic field by characterizing its Fourier spectral density, both in spatial and temporal domains. By measuring power-law spectral exponents $\\alpha$ (spatial) and $\\beta$ (temporal), we studied the photospheric plasma dynamics and its possible implications for the photosphere-corona coupling. In this investigation we used high spatial resolution, high cadence LOS magnetograms and maps of coronal emission from SDO/HMI and SDO/AIA, correspondingly. The utilized data represent NOAA AR 11158.  We determined the spatio-temporal scaling in two stages. First, the scaling was studied by measuring $\\alpha$ for the spatial power spectral density. In the second stage we carried out the temporal analysis, in which we determined temporal scaling of time series corresponding to the net LOS magnetic field. In both stages, average and time-dependent values were measured.  Time-averaged and time-dependent values of the spatial scaling exponent were measured for the turbulence inertial range of scales, which was determined here to be the scales with linear sizes $l=k^{-1}\\approx$ 2 - 20 Mm. Average power spectral density displayed a power law $\\overline{E}(k)\\sim k^{-\\overline{\\alpha}}$ with scaling exponent $\\overline{\\alpha}\\approx 2$. In addition, the time evolution of the power-law exponent shows values greater than 5/3 during the stable phase of the AR, in agreement with \\cite{2010ApJ...720..717A} for flaring ARs. On the other hand, the power law $k^{-2}$ seems to be a characteristic spectrum for MHD turbulence in which the presence of dynamically uniform strong magnetic field favors kinetic and magnetic fluctuations along the field.  Temporal spectral analysis of the data showed that the time series of net signed LOS magnetic flux density displays a power-law spectrum which can be approximated by the Kolmogorov exponent $E(f)\\propto f^{-5/3}$ for an inertial range of temporal scales from several minutes to several days. Time series presenting a power-law spectra with $\\beta\\neq$ 2 are described in terms of the fBm model. In particular, for a fBm time series, $\\beta\\approx 5/3$ implies an anti-persistent behavior with weakly-anticorrelated increments. In the context of photospheric magnetic field evolution this fBm behavior could be an indication of the system seeking for a balance between injection and dissipation of the photospheric magnetic flux. We believe this balance is a signature of fully-developed turbulence that controls the photospheric magnetic flux dynamics and it is present in order to maintain the statistically steady state of the global photospheric magnetic network (\\cite{2001ApJ...561..427S}).  Position-dependent average values of temporal scaling exponent ($\\overline{\\beta}$) indicate that regions with high average magnetic flux densities are typically associated with higher spectral slopes. Exponents $\\overline{\\beta}$ were measured for a range of frequencies corresponding to the spatial inertial range and consistent with Taylor's approximation. We found that only 1 - 3 \\% of the studied image area satisfies the condition $\\overline{\\beta}\\approx\\overline{\\alpha}$. This implies that Taylor's frozen-in-flow turbulence hypothesis is invalid for most of the field of view, including the AR. Consequently, a linear mapping between spatial and temporal behavior using Taylor's hypothesis seems questionable, and a full spatio-temporal characterization of the photospheric magnetic field is required for a complete description of the system turbulent dynamics. We have taken the first initial steps towards such spatio-temporal characterization in this work.  Short-term evolution (minutes to a few hours) of spatial scaling exponent $\\alpha(t)$ captures systematic changes in the spatial distribution of LOS photospheric magnetic field associated with flaring activity. Flare-related changes manifest themselves in both transient (9 - 12 min) and persistent ($\\approx$ 1 h or longer) variations of $\\alpha(t)$ at the time of the flare and immediately afterwards, respectively. Transient variations in $\\alpha(t)$ are most likely associated with artifacts in the magnetic field data, while persistent changes suggest a change in the state of the photospheric field. NOAA AR 11158 produced six M-class flares during its passage through the field of view of the instrument. We detected such systematic changes for two of these flares, in addition to the X-class flare. Although our results support the idea of a back reaction from the corona to the photosphere right after a flare occurs, careful analysis of a larger data set is required to confirm what types of flares are capable of influencing the post-flare state of the photospheric magnetic field.  Careful spatio-temporal analysis of high-resolution photospheric and coronal images such as the one conducted in this study can improve our understating of the physics of solar ARs and the links between the photosphere and corona during flaring activity. Furthermore, photospheric parameters such as the scaling exponents $\\alpha$ and $\\beta$ may also contain advanced information about the coronal flaring activity. We will expand our studies initiated in this paper by inclusion of new ARs into the analysis and by using full vector photospheric magnetic field data. Our ultimate goal is to better understand the physical properties of flaring ARs and to seek for new precursors for pending major solar eruptions."
    },
    "1402/1402.5333_arXiv.txt": {
        "abstract": "Understanding the phases of water ice that were present in the solar nebula has implications for understanding cometary and planetary compositions as well as internal evolution of these bodies.  Here we show that amorphous ice formed more readily than previously recognized, with formation at temperatures $<$70 K being possible under protoplanetary disk conditions.  We further argue that photodesorption and freeze-out of water molecules near the surface layers of the solar nebula would have provided the conditions needed for amorphous ice to form.  This processing would be a natural consequence of ice dynamics, and would allow for the trapping of noble gases and other volatiles in water ice in the outer solar nebula. ",
        "introduction": "Whether water  was present in the outer solar nebula as amorphous or crystalline ice remains a topic of debate.  If amorphous ice formed in the nebula, then it would have been capable of trapping other gases present and could possibly explain the noble gas contents of comets and Jupiter \\citep{barnun85,owen99,yokochi12}.    However, \\citet{kouchi94} demonstrated that as water ice is expected to condense at temperatures of 120-180 K in the solar nebula, water molecules would have been able to arrange themselves into a crystalline structure.  Because the D/H ratios of water throughout the solar nebula is significantly less than that found in the interstellar medium, it is thought that much of the ice in the nebula saw high temperatures near the young Sun then was carried outward to cooler regions, condensing at these temperatures along the way \\citep{mousis00,yang13}.  For example, bringing the D/H ratio of water from the interstellar ratio of $\\sim$0.001-0.01 to the cometary values of 1.5-3$\\times$10$^{-4}$ \\citep{hartogh11} requires that that 70-97\\% of water equilibrated with molecular hydrogen (solar D/H $\\sim$10$^{-5}$) at high temperatures ($>$500 K) in the solar nebula, before cooling and condensing as ice.  This has led many to argue that amorphous ice formation was precluded in the solar nebula, and that all water ice was crystalline requiring another means of trapping volatiles such as clathrates \\citep{kellerjorda01,hersant04}.  The solar nebula is expected to have been a dynamic object, in which gas and dust were pushed around as mass and angular momentum were transported throughout the disk \\citep{armitage11}.  As a consequence of this evolution, gas and dust would be transported throughout the disk and exposed to a variety of physical and chemical environments.  Upon passing through such environments, the dust and gas could be altered by chemical reactions, irradiation, or phase transitions.  This was demonstrated in \\citet{cieslasandford12}, as icy grains in the outer solar nebula were found to be lofted to high altitudes above the disk midplane and exposed to intense radiation fluxes which would break molecular bonds in the ices and allow the resulting ions and radicals to react to form more complex species.  Thus just because a compound or phase is formed in a protoplanetary disk, it does not mean that original phase is preserved throughout the lifetime of the solar nebula.  In this Letter we revisit the conditions under which the various phases of water ice would form in the solar nebula.  In particular, the conditions for amorphous ice formation in the solar nebula are derived assuming that the ice forms on a substrate other than hexagonally crystalline ice.  We then show that the conditions for amorphous ice formation would have been met as icy grains were irradiated by UV photons at the surface layers of the solar nebula and the resulting photodesorbed water froze-out again at greater depths inside the protoplanetary disk. ",
        "conclusions": "The results presented here show that even if water ice condensed in the solar nebula at temperatures in the 120-180 K range, where crystalline ice is expected to form, the dynamical evolution of the icy grains in the outer solar nebula would lead this ice to be lost and reformed as amorphous ice.  Formation of amorphous ice in this manner may allow for the trapping of other gaseous species, such as N- and C- bearing species as well as noble gases.  The trapping of such species in ices would allow volatile-rich planetesimals to form in the outer nebula, and may explain the enhanced abundances of these elements relative to hydrogen in the Jovian atmosphere compared to a gas of solar composition \\citep{owen99,atreya03}.  An issue that requires future attention is the detailed fate of the liberated water at the upper layers of the protoplanetary disk.  If water molecules are desorbed sufficiently high in the disk, they may be photodissociated by the same UV photons that liberate them from the grains they are on.  The extent to which materials are photodissociated will depend on the amount of time materials spend in the extreme upper layers of the disk.  As shown here, the amount of water surrounding a grain will constantly be changing, with molecules freezing out and then being liberated over and over again as the conditions they exposed to are constantly changing.  Thus there will be a time component to determining the precise fraction of water molecules that are liberated around particular grains that gets photodissociated compared to the bulk water present in a given region.  However, estimates of the amount of water vapor present in irradiated regions of the disk under steady-state conditions with no dynamic transport give $n_{H_{2}O}$/$n_{H} \\sim$2$\\times$10$^{-7}$, or 1/1000 the maximum value reached here \\citep{hollenbach09}.  Taking this as an upper limit for the amount of water vapor present in the regions of the disk considered would imply freeze-out fluxes that are ~1000$\\times$ smaller than calculated here, or 10$^{2}$-10$^{6}$ cm$^{-2}$ s$^{-1}$.  These values still exceed the $F_{c}(I_{p})$  (Fig. 1) for temperatures $<$60 K by orders of magnitude, meaning amorphous ice would still form. Further, even with photodissociation, the resulting species will freeze-out again and reform H$_{2}$O on grain surfaces at very rapid rates \\citep{ioppolo08,dulieu10}.  Given the low surface diffusivities of the molecules on $I_{p}$ at the temperatures considered here, when such molecules form they would not be able to reorder themselves, resulting in amorphous ice.  The formation of amorphous ice alone would aid in the continued formation of this phase of water ice: the surface diffusivity of amorphous ice is orders of magnitude smaller than $I_{p}$ \\citep{kouchi94}, meaning any water molecules formed or frozen out on it would remain locked in place and unable to rearrange into a crystalline structure.  This would be true whether the amorphous ice formed via direct freeze-out, surface processes, or through irradiation damage.  \\citet{letobaratta03} found that crystalline water ice would be amorphized after receiving a dose of UV photons of a few eV per molecule, or a few UV photons per molecule.  The dosages calculated here and found by \\citet{cieslasandford12} are beyond those needed for amorphization.  Thus any water ice that forms as a result of photodesorption in the outer solar nebula would likely form on a very rough substrate with limited surface diffusivity, leading the water to freeze out in the amorphous phase.   Should we thus expect all water ice in the solar nebula and protoplanetary disks to become amorphous?  Not necessarily.  Just as we have shown that cycling of this type could lead crystalline ice to be transformed to amorphous ice, amorphous ice could be transformed to crystalline as it migrates through the protoplanetary disk.  At temperatures near $\\sim$100 K, amorphous ice undergoes a spontaneous phase transition, where molecules rearrange themselves to form crystalline ice.  Thus as ice particles migrate radially in a protoplanetary disk and reach higher temperatures, the amorphous ice may be lost.  Trapped volatiles which resided on the grain surface when ice molecules froze out on top of them may be retained, despite the phase transition \\citep{viti04, collings04}, meaning that crystalline ices with trapped volatiles may also be present in such disks.  Determining the relative abundances of each phase of ice requires a detailed investigation of the radial transport of icy grains throughout the evolution of the solar nebula, which should be the goal of future work."
    },
    "1402/1402.2380_arXiv.txt": {
        "abstract": "PSR~J0218$+$4232 is a millisecond pulsar (MSP) with a flux density $\\sim$0.9 mJy at 1.4~GHz. It is very bright in the high-energy X-ray and $\\gamma$-ray domains.  We conducted an astrometric program using the European VLBI Network (EVN) at 1.6~GHz to measure its proper motion and parallax. A model-independent distance would also help constrain its $\\gamma$-ray luminosity. We achieved a detection of signal-to-noise ratio S/N~$>37$ for the weak pulsar in all five epochs. Using an extragalactic radio source lying 20 arcmin away from the pulsar, we estimate the pulsar's proper motion to be $\\mu_{\\alpha}\\cos\\delta=5.35\\pm0.05$~mas\\,yr$^{-1}$ and $\\mu_{\\delta}=-3.74\\pm 0.12$\\,mas\\,yr$^{-1}$, and a parallax of $\\pi=0.16\\pm0.09$\\,mas. The very long baseline interferometry (VLBI) proper motion has significantly improved upon the estimates from long-term pulsar timing observations. The VLBI parallax provides the first model-independent distance constraints: $d=6.3^{+8.0}_{-2.3}$\\,kpc, with a corresponding $3\\sigma$ lower-limit of $d=2.3$\\,kpc. This is the first pulsar trigonometric parallax measurement based solely on EVN observations. Using the derived distance, we believe that PSR~J0218$+$4232 is the most energetic $\\gamma$-ray MSP known to date. The luminosity based on even our 3$\\sigma$ lower-limit distance is high enough to pose challenges to the conventional outer gap and slot gap models. ",
        "introduction": "PSR~J0218$+$4232 is a pulsar with a spin period of 2.3~millisecond and a period derivative of 8.0$\\times$10$^{-20}$~s\\,s$^{-1}$. This millisecond pulsar (MSP) was first discovered by \\citet{navarro95} using the Lovell telescope. The radio timing observations also found that it has a low-mass companion ($M\\gtrsim$~0.16~$M_\\odot$) with an orbital period of 2 days. Optical observations using the Keck telescope revealed that the companion is a helium-core white dwarf with a temperature of $T_\\mathrm{eff}=8060\\pm150$~K, and a distance constraint of 2.5--4~kpc is given by white-dwarf modeling \\citep{bassa03}. Note that the distance uncertainty derived by this method is difficult to quantify exactly, because it is dependent on the white-dwarf's mass and optical luminosity, which are both correlated with the cooling age that is highly uncertain for observations and theoretical white-dwarf models \\citep{bassa03}.  PSR~J0218$+$4232 is also an energetic pulse emitter in X-rays and $\\gamma$-rays. It was a $\\gamma$-ray MSP candidate detected by Energetic Gamma Ray Experiment Telescope \\citep{kuiper00}. Soon after \\textit{Fermi Gamma-ray Space Telescope} was launched, its was confirmed as a $\\gamma$-ray MSP \\citep{msp-sci}. The X-ray pulsed emission has also been well detected and monitored by many X-ray telescopes \\citep[][]{webb04}. Currently, there are approximately 150 rotation-powered pulsars detected in the X-ray band, and about 50 of those have millisecond spin periods \\citep{becker09}.  PSR~J0218$+$4232 is one of a few pulsars with a flux $>10^{-5}$ photon cm$^{-2}$\\,s$^{-1}$ in the 2~--~10~keV band \\citep[e.g.,][]{kuiper02}. An absolute X-ray timing accuracy of $\\sim$200~$\\mu$s was achieved by \\emph{Chandra} \\citep{kuiper04} and 40 $\\mu$s by \\emph{XMM-Newton} \\citep{webb04}.   Astrometric parameters (e.g. position, proper motion, parallax) can be determined from pulsar timing observations over a time span of several years. It is still a challenge to measure the times of arrival of its pulses with a high precision since PSR~J0218$+$4232 has broad profile and significant ($\\sim$50\\%) non-pulsed emission \\citep{navarro95}. The best timing solution published to date was derived from observations at Effelsberg \\citep{laz09}. The proper motion derived from these observations is $\\mu_{\\alpha}\\cos\\delta=+5.1\\pm0.3$~mas\\,yr$^{-1}$ and $\\mu_{\\delta}=-2.3\\pm0.7$\\,mas\\,yr$^{-1}$, whose uncertainties are better than the previously published values in \\citet{hobbs05}. Furthermore, due to the occurrence of sources of timing noise in the time-of-arrival data (e.g., the pulsar's intrinsic spin-down noise, noise induced by the stochastic gravitational wave background, dispersion measure (DM) variations) as well as a relatively short timespan of timing observations that can lead to covariances with other parameters of the timing model, the pulsar-timing method can lead to significant errors in the estimated astrometric parameters.  The proper motion and trigonometric parallax of a pulsar can also be independently measured with VLBI observations. The high-precision VLBI astrometry has been applied to many bright slow pulsars \\citep[e.g.,][]{cam96, brisken02, brisken03, chatterjee09}, and MSPs, such as PSR~B1937$+$21 \\citep{1937}, PSR~J0437$-$4715 \\citep{0437}, and PSR~B1257$+$12 with three planets \\citep{yan13}. A pulsar distance measured to 0.4\\% accuracy has been recently achieved in the VLBI astrometry of PSR~J2222$-$0137 \\citep{del13}.  Astrometric parameters derived from VLBI observations can further improve the estimation of parameters from long-term pulsar timing observations, by providing a prior constraint on astrometric parameters to which the timing analysis is insensitive, but which may themselves be highly covariant with other parameters uniquely approachable via timing.  Combining VLBI- and timing-derived astrometry can contribute to frame ties between the International Celestial Reference Frame and the dynamical solar-system frame, which underlie VLBI and pulsar timing, respectively \\citep{madison13}.  \\begin{figure*}[!tb] \\centering \\includegraphics[angle=-90,width=0.98\\textwidth]{f1.eps} \\caption{EVN astrometry: PSR J0218+4232 positions, $1\\sigma$ error ellipses, and the modeled track in the sky plane. A star symbol denotes the calculated pulsar position at the epoch of each observing session. The origin gives the position derived at the reference epoch J2011.8793. Some results from the least-squares fit are annotated in each panel. Table~\\ref{tbl_2} lists the results for the fit that includes parallax.  } \\label{fig1} \\end{figure*}   In this Letter, we present the results of the first VLBI observations of PSR~J0218$+$4232. We summarize the strategy of the VLBI observations and the post-correlation data reduction in Section 2. We discuss the estimation of the pulsar's astrometric parameters in Section 3. Finally, we address the model constraints on the $\\gamma$-ray luminosity of the pulsar in Section 4. ",
        "conclusions": ""
    },
    "1402/1402.7264_arXiv.txt": {
        "abstract": "We present ages and masses for 601 star clusters in M31 from the analysis of the six filter integrated light measurements from near ultraviolet to near infrared wavelengths, made as part of the Panchromatic Hubble Andromeda Treasury (PHAT). We derive the ages and masses using a probabilistic technique, which accounts for the effects of stochastic sampling of the stellar initial mass function. Tests on synthetic data show that this method, in conjunction with the exquisite sensitivity of the PHAT observations and their broad wavelength baseline, provides robust age and mass recovery for clusters ranging from $\\sim10^2 - 2\\times10^6\\msun$. We find that the cluster age distribution is consistent with being uniform over the past $100 \\Myr$, { which suggests a weak effect of cluster disruption within M31}.  The age distribution of older ($>100\\Myr$) clusters fall towards old ages, consistent with a power-law decline of index $-1$, { likely from a combination of fading and disruption of the clusters}.  We find that the mass distribution of the whole sample can be well-described by a single power-law with a spectral index of $-1.9 \\pm 0.1$ over the range of $10^3-3\\times10^5\\msun$.  However, if we subdivide the sample by galactocentric radius, we find that the age distributions remain unchanged.  However, the mass spectral index varies significantly, showing best fit values between $-2.2$ and $-1.8$, with the shallower slope in the highest star formation intensity regions. We explore the robustness of our study to potential systematics and conclude that the cluster mass function may vary with respect to environment. ",
        "introduction": "It has become clear that a significant fraction of star formation occurs in stellar clusters. However, deriving galaxy histories from observations of clusters is complicated by significant uncertainties.  Controversial questions have been raised regarding cluster properties in various environments as different analyses could lead to different conclusions.  Claims exists for interesting variations or trends in cluster colors, their lifetimes as gravitationally bound objects, and age or mass distributions within or between galaxies, and the evolution from initial to current cluster mass functions \\citep[\\eg,][]{Zepf1993, Kumai1993, Girardi1995, Elmegreen1997, Bastian2011}.  Such studies rely on our ability to estimate intrinsic properties of stellar clusters, in particular, their ages and masses. Much observational effort has therefore been invested in determining the distributions of star cluster ages and masses \\citep{Searle1980, Larsen2000, Billett2002, Hunter2003, Fall2005, Rafelski2005, Dowell2008, Larsen2009, Chandar2010, Bastian2012}, and using the resulting data to determine the dominant mechanisms of cluster formation and disruption \\citep{Kroupa2002,Boutloukos2003, Lamers2005, Whitmore2007, Parmentier2008, Fall2009, Elmegreen2010, Converse2011}.  However, most of the existing work deals with observations of relatively massive clusters (a few $10^4-10^5\\msun$), which are the least affected by disruption processes and thus are the most stable in various environments.  As a result, divergences between disruption models \\citep[\\eg][among many others]{Baumgardt2003, Fall2009} are still subject to debate in the literature \\citep[\\eg][most recently in M83]{Whitmore2011, SilvaVilla2014arXiv}. Only a few studies have been able to probe the smaller clusters that are the most sensitive to environmental effects (mainly in the Galaxy, \\eg, \\citealt{Borissova2011}, or the Magellanic Clouds, \\eg, \\citealt{Popescu2012}).  The Panchromatic Hubble Andromeda Treasury (PHAT; \\citealt{Dalcanton2012a}) is an ongoing multi-cycle Hubble Space Telescope (HST) program that is ideal for studying stellar clusters in M31.  The survey imaged one-third of the M31 disk at high spatial resolution with wavelength coverage from the ultraviolet through the near-infrared.  The sensitivity of the latest HST instruments allow us to detect clusters in M31 down to a regime in which cluster luminosities overlap those of individual bright stars, and hence to very low masses \\citep{Johnson2012a}.  This survey spans a wide range of environments in both star formation intensities and gas densities. This diversity is an advantage for addressing how environment affects cluster formation.  The PHAT survey has already significantly increased the number of clusters known in M31. \\citet{Johnson2012a} identified $601$ stellar clusters using the first quarter of the total PHAT coverage. This new cluster catalog contains more than a factor of four increase in the number of known clusters within the survey area.  Moreover, the uniform photometric coverage from the UV to near-IR allows accurate age-dating of the clusters.  Even this preliminary sample breaks new ground for studying clusters outside the Milky Way and the Magellanic Clouds, probing  about two orders of magnitudes fainter in the luminosity function \\citep[their figure 11]{Johnson2012a}.  This paper is part of a series utilizing the PHAT dataset for studies of stellar clusters.  \\citet{Johnson2012a} presented the first installment of a HST-based cluster catalog, which serves as the basis for an extensive study of Andromeda's cluster population.  In this paper, we focus on the determination of ages and masses of the first year sample, looking forward to the final product of this four year Treasury program. Our estimates of the properties of the clusters are derived from integrated photometry in six broad bands and we especially focus our attention on the characterization of the lowest-mass clusters.  Additional studies, including analysis of structural parameters, resolved star content, and integrated spectroscopy of the cluster sample will follow in subsequent work.  \\medskip  This paper is organized as follows.  \\S \\ref{sec:data} presents the cluster sample and the key elements of their photometry. \\S \\ref{sec:Analysis} describes the analysis and the cluster models used to derive the properties of the clusters, and briefly highlights the possible artifacts of the method using synthetic data.  \\S \\ref{sec:yr1properties} describes our results for the entire sample and for individual regions across M31.  Finally, we discuss those results in \\S \\ref{sec:discussion} before drawing our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have derived ages, masses and exintction for the Year 1 PHAT cluster sample \\citep{Johnson2012a}, by comparing the cluster integrated 6-filter fluxes with an extended version of the stochastically sampled model clusters presented in \\citet{Fouesneau2010}. The locus of the collection of stochastic models in color space (\\eg, Figure 2) shows excellent agreement with that of the collection of cluster observations. Clusters with broadband colors either bluer or redder than those of the traditional continuous models find a natural match with the models we used in this paper.  We generated the full joint probability distribution function of the age, mass, and extinction for each of the 601 individual clusters in the sample. We then combined their individual distributions into global cluster age and mass distributions, noting limits at which completeness issues in the sample become severe.  The sample of clusters spans the entire length of the age sequence and includes a significant number of clusters with masses well below $10^3\\msun$. Only a few datasets have the ability to sample objects across a variety of stages in cluster evolution over such a large, uninterrupted mass range.  We find that the cluster age distribution shows a constant number of clusters over the last $\\sim 100\\Myr$, with a power-law decline at older ages (see Figs\\,\\ref{fig:marginal_age_dist} \\& \\ref{fig:dndt_dist_per_brick}).  At least above the mass of $10^{3.2}\\msun$, these results are consistent with M31 producing a constant number of clusters from $100\\Myr$ ago to present, with little significant cluster disruption over this timescale.  The mass distribution derived from the analysis closely resembles the power-law distributions obtained from many other galaxies.  Specifically, the overall power-law index of the mass distribution is consistent with the canonical value of $-$2. However, the current cluster sample suggests a possible radial variation of this distribution across the disk, with the shallowest power-law found in the region with the highest star formation rate.  When we study the entire PHAT survey, including lower masses and a larger sample of fainter clusters, the improved accuracy and time resolution achievable with the new stochastic methods will allow us to address new questions. Future work will account for the challenging determination of completeness and selection effects.  In particular, the expected number of clusters in PHAT will eventually provide 5 times more clusters over a broad range of local environments, which will open the possibility to study local variations among cluster populations beyond our current the initial assessment in this study."
    },
    "1402/1402.2488_arXiv.txt": {
        "abstract": "{The structure of protoplanetary disks is thought to be linked to the temperature and chemistry of their dust and gas. Whether the disk is flat or flaring depends on the amount of radiation that it absorbs at a given radius, and on the efficiency with which this is converted into thermal energy. The understanding of these heating and cooling processes is crucial to provide a reliable disk structure for the interpretation of dust continuum emission and gas line fluxes. Especially in the upper layers of the disk, where gas and dust are thermally decoupled, the infrared line emission is strictly related to the gas heating/cooling processes.}{We aim to study the thermal properties of the disk in the oxygen line emission region, and to investigate the relative importance of X-ray (1-120 \\AA) and far-UV  radiation (FUV, 912-2070 \\AA) for the heating balance there.}{We use \\on\\,63 \\mic\\,line fluxes observed in a sample of protoplanetary disks of the Taurus/Auriga star forming region and compare it to the model predictions presented in our previous work. The data were obtained with the PACS instrument on board the Herschel Space Observatory as part of the Herschel Open Time Key Program GASPS (GAS in Protoplanetary diskS), published in Howard et al. (2013).}{Our theoretical grid of disk models can reproduce the \\on\\,absolute fluxes and predict a correlation between \\on\\,and the sum \\lx+\\luv. The data show no correlation between the \\on\\,line flux and the X-ray luminosity, the FUV luminosity or their sum.}{The data show that the FUV or X-ray radiation has no notable impact on the region where the \\on\\,line is formed. This is in contrast with what is predicted from our models. Possible explanations are that the disks in Taurus are less flaring than the hydrostatic models predict, and/or that other disk structure aspects that were left unchanged in our models are important. Disk models should include flat geometries, varying parameters such as outer radius, dust settling, and the dust-to-gas mass ratio, which might play an equally important role for the \\on\\,emission. To improve statistics and draw more robust conclusions on the thermal processes that dominate the atmosphere of protoplanetary disks surrounding T\\,Tauri stars, more \\luv\\,and \\lx\\,measurements are needed. High spatial and spectra resolution data is required to disentangle the fraction of \\on\\,flux emitted by the disk in outflow sources.} ",
        "introduction": "Planet formation is strongly linked to the physical properties of the parent disk. Important constraints on the timescale for the gas accretion of giant planets are posed by photoevaporation models. The results of such models are essential in order to estimate the mass loss rates, and hence the survival time of gas in disks (\\citealt{Ale06,Erc08,GDH09}). The stellar radiation, especially in the high energy regime (E $>$ 6 eV), is responsible for the thermo-chemical conditions in the disk atmosphere, as it provides most of the energy that causes the gas temperature to exceed the dust temperature there \\citep{Kam04,Jon04,Gla04}. However, the thermal processes that heat and shape protoplanetary disks are poorly constrained and can only be indirectly measured through cooling lines. One of the dominant cooling lines that can be used to understand these processes is the 63 micron line of neutral oxygen \\citep{Gor08,Mei08,Woi09,Are12}.  T\\,Tauri stars emit radiation at high energies, due both to chromospheric activity and accretion of disk material onto the stellar surface. The FUV luminosity between 7 and 10 eV ($\\Delta\\lambda$=1240-1770 \\AA), has been measured by \\citet{Yan12} for a sample of accreting sources in Taurus: they found values between 10$^{30}$ and few times 10$^{32}$ erg/s. The emission is in excess when compared to the stellar emission in the same energy band for non-accreting young stars of the same spectral type. This suggests that accretion is responsible for this emission, in which case it is caused by shocks created by the magnetic field that channels disk material toward the stellar surface \\citep{Cal98,Val00}. EUV ($\\Delta\\lambda$= 120-912 \\AA, $\\Delta$E=13.6-100 eV) radiation is believed to mainly affect the upper disk surface at small radii, as the high cross section for absorption only allows penetration of small columns of $N_{\\rm H}\\sim 10^{19}$ cm$^{-2}$. The XEST survey \\citep{Gue07} has shown that young stars are also active X-ray emitters, mainly due to chromospheric activity, and can reach luminosities between 10$^{29}$ and 10$^{31}$ erg/s. The high energy depositions ($>$0.01 $L_{*}$) and heating efficiencies ($\\sim$ 30\\%) of X-rays cause the tenuous disk atmosphere to heat up to temperatures of the order of a few thousand Kelvin \\citep{Gla07,Nom07,Gor08,Erc08,Are11}.  Recent observations, carried out with the Herschel Space Observatory toward the Taurus forming region, offer the chance to test model predictions on the thermal structure of the region where the \\on\\,63.2 \\mic\\,line is emitted. This line is predicted to arise from the disk atmosphere in the radial region between a few 10 AU and 200 AU \\citep{Woi09,Are12}. The emission region is directly exposed to the stellar radiation and models suggest that FUV and X-ray radiation are the main heating agents there. PAH and dust photoelectric heating as well as Coulomb heating, cause the gas temperature to be of the order of $\\sim$ 200-300 K \\citep{Gor08,Mei08,Mei12}.  In this work, we explore possible correlations of the \\on\\,emission with X-ray luminosity and FUV luminosity, and compare the \\citet{Are12} model predictions for the \\on\\,63 \\mic\\,emission with data collected within the GASPS (GAS in Protoplanetary DiskS, P.I. Dent) Open Time Key Program, taken with the PACS instrument on board the Herschel Space Observatory \\citep{Den13}.  In the following, we make the hypothesis that most of the \\on\\,emission is produced in the disk. Outflow sources, that have on average higher accretion rates, will then produce more FUV radiation and thus stronger FUV illumination of the disk surface and stronger line emission. %  In Sect.\\ref{obser} we present the collected observational data set and in Sect.\\ref{models} we explain the main findings of the models studied in \\citet{Are12}. In Sect.\\ref{results} we show the results of the comparison between model predictions and observations, these will be discussed in Sect.\\ref{disc}. Conclusions and remarks about future work are summarised in Sect.\\ref{fut}. ",
        "conclusions": "\\label{fut} In this work we studied the impact of FUV and X-ray radiation on the thermal balance in the oxygen emission region for protoplanetary disks surrounding T\\,Tauri stars. We compared disk model predictions with observations of the \\on\\,63\\mic\\,line toward protoplanetary disks that do not show outflow emission in the Taurus region obtained with the PACS instrument on board Herschel.  The observations show no correlation between the \\on\\,63 \\mic\\,line emission and the X-ray luminosity or the FUV luminosity or with their sum.  Our thermo-chemical disk models calculated with ProDiMo, show that our predictions on the \\on\\,fluxes qualitatively agree with the observations. There is no correlation between \\on\\,and \\lx\\,or \\luv, as the data suggest. Nevertheless, the models predict a correlation between \\on\\,and the sum \\lx+\\luv, which is not seen in the data. The reason can be the limited set of parameters varied in our model (\\lx, \\luv, grain minimum size, power law of the grain size distribution and power law of the surface density distribution) grid to understand the relative importance of \\lx\\,and \\luv.However, other parameters can affect the \\on\\,line, causing the correlation we predict to vanish when a more complete grid is used.  To include all the disk parameters that influence the \\on\\,line, a different set of models should be used. Flatter disk geometries should be included as well as a proper treatment of X-ray and FUV physics and dust settling (local variations of the gas-to-dust mass ratio).  Moreover, such models should be compared to a higher number of observations: high spatial and spectral resolution data is required to disentangle the location of the emission region of the line. In many sources that drive outflows, the contribution of the disk to the total flux of the line remains unclear.  More measurements of \\luv\\,would also be necessary. To test the threshold mechanisms proposed in our previous work, observations of \\on\\,of sources with \\lx $>$ \\luv, if any, are essential.  The understanding of the [OI] dependence on the FUV and X-ray radiation gives the possibility to investigate the gas surface layers above the H/H2 transition. Such studies are very interesting for understanding the photoevaporation mechanism and how it may drive disk evolution across the transition from optically thick to debris disk. \\begin{figure}[h!] \\centering \\includegraphics[scale=0.4,angle=-90]{fit03.pdf}% \\caption{\\small \\on\\,flux versus \\lx+\\luv. Model points are black while data points are red. In our models we considered a sun-like star surrounded by a disk of 0.02 $M_{\\odot}$ which extends from 0.5 to 500 AU.} \\label{fit3} \\end{figure}  \\tiny \\emph{Acknowledgements}.  We thank Aki Roberge for her comments on the FUV analysis which helped to improved the paper and Glenn White for thoroughly reading the paper. WFT, PW, FM, MG and IK acknowledge funding from the EU FP7-2011 under Grant Agreement nr. 284405. L.P. acknowledges the funding from the FP7 Intra-European Marie Curie Fellowship (PIEF-GA-2009-253896)."
    },
    "1402/1402.1507.txt": {
        "abstract": " ",
        "introduction": "One of the most exciting open questions at the interface between particle physics and cosmology is the nature of Dark Matter (DM). The first person who provided evidence and inferred the presence of DM was a Swiss-American astrophysicist, Fritz Zwicky. He applied the virial theorem to the Coma cluster of galaxies and obtained evidence of unseen mass. Roughly 40 years following the discoveries of Zwicky and others, Vera Rubin and collaborators conducted an extensive study of the rotation curves of isolated spiral galaxies. They announced the discovery that the rotational curves of stars in spiral galaxies exhibit a characteristic flat behaviour at large distance in contrast with Kepler's law. Many other evidence of unseen mass on distance scales of the size of galaxies and clusters of galaxies appeared throughout the years, but the most precise measurement of the total amount of DM comes from cosmological scales. In particular, the measurements of modern precision cosmology (the Cosmic Microwave Background (CMB) and the surveys of the Large Scale Structure (LSS) of the Universe), provide the current most relevant evidence. Apart from the qualitative agreement, it is the quantitative fitting of the wealth of available data that allows the amount of DM to be one of the cosmological parameters now most precisely measured ($\\Omega_\\chi h^2=0.1199\\pm0.0027$, see Tab.~2 of~\\cite{Ade:2013zuv}).  Therefore we have compelling evidence of unseen mass, but the microscopic features of this new kind of matter remain unknown yet. Direct and Indirect searches may shed light on the nature of DM, and therefore a careful study of their phenomenology is fundamental. For a pedagogical review on this subject, see e.g.~\\cite{Bertone:2004pz}.  \\medskip Direct searches for DM aim at detecting the nuclear recoils arising from scattering between DM particles and target nuclei in underground detectors. DM direct detection experiments are providing exciting results in terms of measured features which have the right properties to be potentially ascribed to a DM signal.  For example in addition to the long-standing \\DAMA\\ results, nowadays there are other experiments, such as \\CoGeNT, \\CRESST\\ and \\CDMSSi\\, that start to see some anomalies in their counting rates. On the other hand ,the situation in this field is extremely unclear and confusing, because on top on these positive result experiments, the constraints coming from null results, like \\XENONhundred, \\COUPP, \\PICASSO\\ and very recently \\LUX, are very stringent and put  the interpretation of the anomalies in terms of a DM interaction in serious trouble. Nevertheless, there are at least two main caveats when the results from the experiments commented upon above are interpreted. The first is that one has to treat with great care the fine experimental details associated with the results quoted by each experiment. The second caveat is instead associated with the interpretation of the data within a very simple-minded DM model. For instance, the DM-nucleus spin independent contact interaction is just a benchmark example. Upon relaxing some of these assumptions, the current complicated experimental puzzle can probably be solved.  \\medskip The {\\bf scope of this work} is to present the {\\em status} of direct DM detection with specific attention to the experimental results and their phenomenological interpretation in terms of DM interaction. In particular in Sec.~\\ref{Basics} I review a new and more general approach to study signals in this field based on non-relativistic operators which parameterize most efficiently the DM-nucleus interactions in terms of a very limited number of relevant degrees of freedom. In Sec.~\\ref{DDStatus} I review the experimental results and their interpretation in terms of the ``standard'' spin independent (SI) interaction. I list then the main uncertainties that affect the theoretical interpretation of the data: this is a very promising area of research since only major advancements here can probably reconcile the complicated puzzle showed  by the experiments up. Finally in Sec.~\\ref{LRInteraction}, I pose my attention on the uncertainties coming from the nature of the interaction. In particular the phenomenology  of a class of models in which the interaction between DM particles and target nuclei is of a long range type is discussed. ",
        "conclusions": "Direct DM searches is now characterized by tantalizing results and hints   that make this field very active both from the theoretical and experimental side. In particular, in addition to the long standing \\DAMA\\ results, nowadays there are other experiments, like \\CoGeNT, \\CRESST\\ and \\CDMSSi\\ that are starting to observe anomalies in their counting rates. On the other hand, the increasingly stringent constraints coming from null result experiments put in serious trouble the theoretical interpretation of the data, at least in terms of the simple-minded SI contact interaction. In this work I  discussed the {\\em status} of direct DM detection with specific attention to the experimental results and their phenomenological interpretation in terms of DM interaction. In particular, in the first part I presented a new and more general approach to study signals in this field based on non-relativistic operators. Then I reviewed  the experimental results and their interpretation in terms of the ``standard'' SI interaction pointing out all the uncertainties which enter in this field. In the last part of this work, I investigated a fermionic Dark Matter particle carrying a small milli-charged and analyzed its impact on direct detection experiments. I showed that this kind of long range interaction can accommodate the positive experimental results. By assuming a conservative choice for the lower threshold of the \\XENONhundred\\ and \\LUX\\ experiments I have demonstrated that this candidate is not ruled out. I also determined the complementary class of constraints which are relevant for milli-charged DM particles with long range forces.  \\medskip Finally,  I would like to propose a possible direction to pursue  in order to make sense of the current exciting experimental panorama based on the formalism of non-relativistic operators. Indeed, as we have seen in the first part of this work, it allows us to describe the DM-nucleus interactions in terms of a very limited number of relevant degrees of freedom. In this way, it is possible to parametrize the model-dependent part of the rate from the model-independent one encapsulated in a sort of \\emph{integrated form factors} that encode all the dependences on the astrophysics, nuclear physics and experimental details. Therefore, since one is ignorant or agnostic about the underlying theory, I would like to encourage a synergy between nuclear physicists and experimentalists in order to provide a complete set of \\emph{integrated form factors} defined in Eq.~\\eqref{integratedFF}. It would be extremely useful for the community, because in this way, one can compute the expected number of events for any kind of interaction (e.g.~including isospin-violating interactions, momentum-dependent form factors, velocity-dependent form factors) and compare  directly with the experimental results. Providing the \\emph{integrated form factors} will thus be  the first step towards a model-independent analysis in direct DM searches.   \\small \\paragraph{Acknowledgements} We thank Marco Cirelli, Eugenio Del Nobile and Joe Silk for useful discussions.  \\bigskip \\noindent The author declare that there is no conflict of interests regarding the publication of this article.  \\bigskip \\appendix  \\footnotesize \\begin{multicols}{2}"
    },
    "1402/1402.0194.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) {Galactic winds are a common phenomenon in starburst galaxies in the local universe as well as at higher redshifts. Their sources are superbubbles driven by sequential supernova explosions in star forming regions, which carve out large holes in the interstellar medium and eject hot, metal enriched gas into the halo and to the galactic neighborhood.    } % {} leave it empty if necessary %{To investigate the physical nature of the `nuc\\-leated instability' of %proto giant planets, the stability of layers %in static, radiative gas spheres is analysed on the basis of Baker's %standard one-zone model.} % aims heading (mandatory) {We investigate the evolution of superbubbles in exponentially stratified disks. We present advanced analytical models for the expansion of such bubbles and calculate their evolution in space and time. With these models one can derive the energy input that is needed for blow-out of superbubbles into the halo and derive the break-up of the shell, since Rayleigh-Taylor instabilities develop soon after a bubble starts to accelerate into the halo. } % In our paper, we analyze the behavior of blow-out superbubbles. % methods heading (mandatory) {The approximation of Kompaneets is modified in order to calculate velocity and acceleration of a bubble analytically. Our new model differs from earlier ones, because it presents for the first time an analytical calculation for the expansion of superbubbles in an exponential density distribution driven by a time-dependent energy input rate. The time-sequence of supernova explosions of OB-stars is modeled using their main sequence lifetime and an initial mass function. } % results heading (mandatory) {We calculate the morphology and kinematics of superbubbles powered by three different kinds of energy input and we derive the energy input required for blow-out as a function of the density and the scale height of the ambient interstellar medium. The Rayleigh-Taylor instability timescale in the shell is calculated in order to estimate when the shell starts to fragment and finally breaks up. Analytical models are a very efficient tool for comparison to observations, like e.g. the Local Bubble and the W4 bubble discussed in this paper, and also give insight into the dynamics of superbubble evolution.} %% %% DB: Satz ge\u00e4ndert %% % % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "% Most massive stars are born in OB-associations in coeval starbursts on timescales of less than 1-2 Myr \\citep{m99}. % These associations can contain a few up to many thousand OB-stars, so-called super star clusters, but typically have 20-40 members \\citep{mk87}. % Energy and mass are injected through strong stellar winds and subsequent supernova (SN) explosions of stars with masses above 8 M$_{\\odot}$. The emerging shock fronts sweep-up the ambient interstellar medium (ISM) and, as the energy input in form of SN-explosions continues, superbubbles (SBs) are produced, which may reach dimensions of kiloparsec-size \\citep[e.g.][]{tsm03}. The swept-up ISM collapses early in the evolution of the SB into a cool, thin, and dense shell \\citep{cmw75} present in HI and H${\\alpha}$ observations. % The bubble interior contains hot ($> 10^6 \\, $K), rarefied material, usually associated with extended diffuse X-ray emission \\citep{sta05}. % % Due to the stratification of the ISM in disk galaxies, the superbubbles can accelerate along the density and pressure gradient and blow out into the halo, appearing as elongated structures. Examples of huge bubbles and supergiant shells are the Cygnus SB with a diameter of 450 pc in the Milky Way \\citep{ccb80} and the Aquila supershell extending at least 550 pc into the Galactic halo \\citep{m96}. %% %% DB: n\u00e4chster Satz ge\u00e4ndert %% Our solar system itself is embedded in an HI cavity with a size of a few hundred parsecs called the Local Bubble (\\citeauthor{l03} \\citeyear{l03}), and is most likely generated by stellar explosions in a nearby moving group (\\citeauthor{bb02} \\citeyear{bb02}). % Also in external systems like in the LMC \\citep{cm90}, in NGC 253 \\citep{s06} and M101 \\citep{ksh91} such bubbles, holes and shells are observed.\\\\ % The acceleration of the shell promotes Rayleigh-Taylor instabilities and after it is fully fragmented, only the walls of the SB are observed. Through such a chimney the hot pressurized SN-ejecta can escape into the halo. The walls may be subject to the gravitational instability and interstellar clouds can form again, which triggers star formation \\citep[e.g.][]{mk87}. This is observed, for example, on the border of the Orion-Eridanus SB \\citep{lc09}.\\\\ % % The knowledge of SB evolution is crucial for understanding the so-called disk-halo connection and it also gives us information about the chemical evolution of the galaxies, the enrichment of the intergalactic as well as the intracluster medium. The thick extraplanar layer of ionized hydrogen seen in many galaxies has probably been blown out of the disk into the halo by photoionization of OB-stars and correlated SNe \\citep{t06}. % With the high star formation rate of starburst galaxies, the energy released by massive bursts of star formation can even push the gas out of the galactic potential, forming a galactic wind. Outflow rates of $0.1-10 \\, $M$_{\\odot}$/yr are common for starburst driven outflows \\citep{bvc07}. % %% %% DB: ich denke im folgenden Satz meinst Du \"metallicity\" anstatt \"abundance\" %% If the hot and metal enriched material is brought to the surrounding intergalactic medium, it will mix after some time, increasing its metallicity. Galactic winds are observed in nearby galaxies \\citep[e.g.][]{dwh98}, as well as up to redshifts of $z \\sim 5$ \\citep[e.g.][]{sw07,dss02}, thus being an ubiquitous phenomenon in star forming galaxies. % If the energy input is not high enough, the gas will fall back onto the disk due to gravity, after loss of pressure support, forming a galactic fountain \\citep{sf76,a00}. % %% %% DB: Satz erweitert %% In this case, the heavy elements released by SN-explosions are returned back to the ISM in the disk, presumably spread over a wider area, and future generations of stars will incorporate them.  Spitoni et al. (2008), also using the Kompaneets approximation, have investigated the expansion of a superbubble, its subsequent fragmentation and also the ballistic motion of the fragments including a drag term, which describes the interaction between a cloud and the halo gas. In addition, these authors have analyzed the chemical enrichment of superbubble shells, and their subsequent fragmentation by Rayleigh-Taylor instabilities, in order to compare the [O/Fe]-ratios of these blobs to high velocity clouds (HVCs). They find that HVCs are not part of the galactic fountain, and even for intermediate velocity clouds (IVCs), which are in the observed range of velocities and heights from the galactic plane, the observed [O/Fe]-ratios can only be reproduced by unrealistically low initial disk abundances. The chemical enrichment of the intergalactic medium will be the subject of a forthcoming paper.\\\\ % This paper is structured as follows: Section 2 shows how the evolution of superbubbles can be described. In Sect.~3 we present the results of this work, which is mainly the expansion of a bubble in space and time and also the onset of Rayleigh-Taylor instabilities in the shell. In Sect.~4 the models are used to analyze two Milky-Way superbubbles. A discussion follows in Sect.~5 and summary and conclusions are presented in Sect.~6. % % % % %________________________________________________SECTION 2_________________________________________________________ % ",
        "conclusions": "% % %% %% DB: hier w\u00fcrde ich KA ausschreiben; viele Leute lesen ja leider nur die Summary %% We have developed analytical models based on the Kompaneets approximation (KA) in order to derive in a fairly simple and straightforward manner the physical parameters of observed superbubbles and their ambient medium and to gain physical insight into the blow-out phenomenon associated with star forming regions.\\\\ % %% %% DB: hier w\u00fcrde ich noch folgenden Satz einf\u00fcgen %% In this paper we have deliberately refrained from building more complex models, which e.g. include stellar wind and Wolf-Rayet wind phases, because we have put our focus on the important dynamical phenomena of blow-out and fragmentation of the outer shell, and their dependence on the energy input source over time. A more detailed description of superbubble evolution models including further stellar evolutionary phases will be the subject of a forthcoming paper.\\\\ % In our work the key aspect was to work out analytically the dynamics of (unfragmented) superbubbles for different energy input modes. We modified the KA to implement a more realistic way of energy input, i.e. modeling the time sequence of exploding stars in an OB association by including the main sequence lifetime of the massive stars and describing the numbers per mass interval by an initial mass function. We tested three different IMF slopes and also compared the IMF-model to a simple SN-model with an instantaneous release of energy and to a wind model with a constant energy input rate. Two different density distributions of the ISM were applied, a symmetric medium with parameters of the Lockman layer and a high-density, low-scale height pure exponential atmosphere with the star cluster dislocated from the galactic plane. Velocity and acceleration of the shock front can be calculated analytically and the question how many SNe are needed for blow-out into a galactic halo can be answered. The exact position of the outer shock in scale height units when the acceleration starts can be given. Furthermore, the timescale for the development of Rayleigh-Taylor instabilities in a SB shell is calculated, and thus, a fragmentation timescale can be derived. The overall pattern shows that at larger scale heights ($H > 400 \\,$pc), independent of the ISM density, the SN-model needs the highest energy input, followed by the wind-model, whereas an IMF with a steep slope is the most efficient one. The same ranking applies to the blow-out timescales of these models (Figs.~\\ref{time_SN}-\\ref{time_wind} $\\&$ Table \\ref{accel}). At low scale heights ($H \\sim 100 \\, $pc) and moderate or high densities ($n_0 \\ge 5 \\,$cm$^{-3}$), the picture changes completely, i.e. the SN-model requires the lowest energy input and the IMF-model with $\\Gamma_3=-1.7$ is the least efficient one. The explanation is that a single release of energy is more powerful in sweeping up a thin layer of ISM, whereas it is easier for the IMF model with an \\emph{increasing} energy input with time ($L_{\\rm{SB}} \\propto t^{\\delta}$, $0 < \\delta < 1$, Eq.~\\ref{eLSB2}) to sustain the supply of energy over a larger distance. % % % When comparing fragmentation timescales (Figs.~\\ref{frag_SN}-\\ref{frag_wind} $\\&$ Table \\ref{frag}), the IMF-models exhibit the lowest values, favouring flatter IMFs with $\\Gamma_1$ and $\\Gamma_2$. In terms of $\\tilde{y}$ and absolute time, fragmentation happens first for the $\\Gamma_2$-model for SBs driven by the minimum number of SNe for blow-out.\\\\ % % Still, the KA is a rather simple model. It does not account for magnetic fields, ambient pressure, inertia and evaporation of the shell. Also a galactic gravitational field and cooling of the shocked gas inside the cavity are neglected. \\\\ % In our model, we include galactic gravity in a rudimentary way: if a bubble fulfills the blow-out criterion and additionally the shell's acceleration at the top of the bubble at the time of fragmentation exceeds the vertical component of the gravitational acceleration in the disk, the SN-ejecta and fragmenting shell are expelled into the halo. We find that this is true for all bubbles created by an energy input of $N_{\\rm{blow}} \\cdot E_{\\rm{SN}} $.\\\\ % Further analysis of a fragmenting superbubble would also involve to calculate the motion of the shell fragments, and ultimately the dynamics of a galactic fountain like e.g. in \\citet{s08}. A simple ballistic treatment of the motion of blob fragments in a gravitational potential would, however, be too simplistic, as they would for some time experience a drag force due to the outflowing hot bubble interior. The drag would be proportional to the ram pressure of the hot gas and the cross section of a blob, which would, to lowest order, be proportional to the thickness of the fragmented shell; and finally it would depend on the geometry of the blob, which might be taken as spherical. However, due to compressibility, bow-shocks and head-tail structures might subsequently be formed, so that for a realistic treatment numerical simulations would be the best choice.\\\\ % According to MM88 cooling of the bubble interior can be neglected for Milky Way type parameters of the ISM, but should be taken into account for smaller OB-associations or a dense and cool ISM. %% %% DB: wieso wolltest Du den folgenden Satz nicht ins Paper nehmen? Er ist auf jeden Fall richtig.  -- ok %% In general, cooling is not important as long as the timescale for radiative cooling is large compared to the characteristic dynamical timescale of the superbubble. % % % It was also shown by MM88 that galactic gravity does not need to be included, because the acceleration of high energy superbubbles exceeds the local gravity. % %. Including a magnetic field would be quite important, but this goes beyond the scope of our analytical model and is left to numerical simulations. \\citet{fmz91} find that the presence of a magnetic field could slow down the expansion of a superbubble. \\citet{s09} argue that the scale height and age of a bubble are underestimated by $\\sim 50 \\, \\%$ when using a Kompaneets model without magnetic fields. However, they cannot produce such narrow superbubbles like W4 with their MHD simulations. Moreover, 3D high resolution numerical simulations \\citep{ab05} show that magnetic tension forces are much less efficient in 3D than in 2D in holding back the expanding bubble.\\\\ % A slightly slower growth of a bubble would be also achieved by taking into account the inertia of the cold massive shell in the calculations. MM88 find a difference of $\\sim 10 \\, \\%$ in radius after comparison with the models of \\citet{s85}, which neglect inertia. Also due to the ambient pressure of the ISM superbubbles should expand more slowly as it was suggested by \\citet{og04}. % In our analytical calculations we find that these effects have to be compensated when we make comparisons to observed bubbles by including a rather high ISM density to prevent SBs from expanding too fast. %% %% DB: was ist damit gemeint? Wieso kann man KA nicht auf \"existierende\" SB anwenden?  - gemeint ist, wenn schon eine fr\u00fchere Generation von Superbubbles in der Umgebung dagewesen ist -- lasse es weg %% Furthermore, a clumpy ambient medium cannot be considered by the KA.\\\\ % % We applied our models to the W4 superbubble and the Local Bubble, both in the Milky Way. It is certainly not easy to compare a simple model with an observed SB, which is not isolated. In the region of the W3/W4/W5 bubbles, several complexes and clouds are found and multiple epochs of star formation make it difficult to distinguish, which cluster has formed which bubble at what time. However, it is most likely that the cluster OCl 352 is responsible for driving the evolution of the bubble \\citep{w07}. \\citet{o05} suggest that winds or SNe of previous stellar generations are responsible for earlier clearing of this region, which could explain the low scale height of around 30 pc. Our calculations suggest that the bubble is younger than found by other authors, which is due to the offset of the association more than one scale height above the Galactic plane. %% %% DB: Satz ge\u00e4ndert, ok? ja %% Shifting to lower densities makes it easier to produce a blow-out superbubble in a shorter timescale. This is an important result and should be included in the models.\\\\ The Local Bubble is one of the rare cases for a double-sided bubble, which can be tested with our symmetric superbubble model. From geometrical properties, we estimate an ISM scale height of $\\sim 80 \\,$pc. In order to reproduce size and age of the bubble correctly, we deduce from our models that it was an intermediate density region ($n_0 \\sim 7 \\,$cm$^{-3}$) before the first SN-explosion around 14 Myrs ago. Also this place in the Milky Way is very complex (neighbouring Loop I superbubble), which can't be included in our modeling of superbubbles. \\\\ We conclude that blow-out energies derived in this paper are lower thresholds and might be higher if e.g. magnetic fields play a role. Accordingly, fitting the models to observed bubbles gives an upper limit for densities of the ambient ISM prior to the first SN-explosion. %% %% DB: Noch ein Schlusswort -- sehr gut, danke %% Observers are encouraged to use the model presented here for deriving important physical parameters of e.g. the energy input sources (number of OB stars, richness of cluster etc.), scale heights, dynamical time scales among other quantities. The solutions of the equations derived in detail here are easy to obtain by simple mathematical programs. Theorists may find it useful to compare our analytic results to high resolution numerical simulations in order to separate more complex effects, such as turbulence, mass loading, magnetic fields etc. from basic physical effects, incorporated in our model. % % % % % % %% %% DB: Du k\u00f6nntest auch noch der armen TU f\u00fcr ein paar Gastaufenthalte danken;-)  -->> nat\u00fcrlich !!! %%"
    },
    "1402/1402.6324_arXiv.txt": {
        "abstract": "We present \\textit{XMM-Newton} and \\textit{Chandra} observations of the low-mass X-ray binary XSS J12270--4859, which experienced a dramatic decline in optical/X-ray brightness at the end of 2012, indicative of the disappearance of its accretion disk. In this new state, the system exhibits previously absent orbital-phase-dependent, large-amplitude X-ray modulations with a decline in flux at superior conjunction. The X-ray emission remains predominantly non-thermal but with an order of magnitude lower mean luminosity and significantly harder spectrum relative to the previous high flux state. This phenomenology is identical to the behavior of the radio millisecond pulsar binary PSR J1023+0038 in the absence of an accretion disk, where the X-ray emission is produced in an intra-binary shock driven by the pulsar wind.  This further demonstrates that XSS J12270--4859 no longer has an accretion disk and has transformed to a full-fledged eclipsing ``redback'' system that hosts an active rotation-powered millisecond pulsar. There is no evidence for diffuse X-ray emission associated with the binary that may arise due to outflows or a wind nebula. An extended source situated 1.5$'$ from XSS J12270--4859 is unlikely to be associated, and is probably a previously uncatalogued galaxy cluster. ",
        "introduction": "Since the original discovery of millisecond pulsars \\citep[MSPs;][]{Back82} the prevailing theory of their formation has centered on ``recycling'' by transfer of matter and angular momentum from a close companion star during a low-mass X-ray binary (LMXB) phase \\citep{Alp82,Rad82}.  The discovery of the first accreting X-ray MSP, SAX J1808.4--3658 \\citep{Wij98}, provided compelling evidence for this evolutionary sequence \\citep[see also][for a review]{Pat12}. Additional observational support in favor of this hypothesis came with the discovery of the ``missing link'' radio pulsar PSR J1023+0038 (also known as FIRST J102347.6+003841). Optical observations revealed an accretion disk in the system in 2001 \\citep{Wang09}, which was absent after 2003 \\citep{Thor05} and at the time of the radio pulsar discovery in 2009 \\citep{Arch09}.  The long-suspected evolutionary connection between LXMBs and radio MSPs was conclusively established when PSR J1824--2452I in the globular cluster M28 was seen to switch between rotation-powered (radio) and high-luminosity accretion-powered (X-ray) pulsations \\citep{Pap13}.  Recent radio, X-ray, and $\\gamma$-ray observations revealed that in 2013 June PSR J1023+0038 had undergone another transformation \\citep{Stap14}. Complete cessation of radio pulsations was observed, accompanied by an extraordinary five-fold increase in the \\textit{Fermi} Large Area Telescope (LAT) flux and enhancement in UV and X-ray brightness \\citep{Pat14,Ten14}.  These recent findings have revealed that rather than a one-time change from a LMXB to a permanent radio MSP state, for some systems this phase of neutron star compact binary evolution involves recurrent switching between the two states.  In the absence of an accretion disk, PSRs J1023+0038 and PSR J1824--2452I belong to the family of so-called ``redbacks'' \\citep[see][and references therein]{Rob11}, namely, eclipsing radio MSPs with relatively massive ($\\gtrsim$0.2 M$_{\\odot}$) non-degenerate companions that are (nearly) Roche-lobe filling. These objects are distinct from ``black widow'' systems, which are bound to very low mass degenerate stars being ablated by the pulsar wind \\citep[][]{Fru88,Stap96}. Extensive X-ray studies of redbacks in globular clusters \\citep{Bog05,Bog10,Bog11a} and the field of the Galaxy \\citep{Arch10,Bog11b,Bog14,Gen13} have revealed that the dominant source of X-rays from these binaries is non-thermal radiation with $\\Gamma\\approx 1-1.5$ originating in an intra-binary shock \\citep{Arons93} generated by the interaction of the pulsar wind and matter from the companion star.  The X-ray emission is strongly modulated at the orbital period in all redback systems, with a decline in flux when the secondary star is between the pulsar and the observer. This orbital-phase-dependent X-ray variability appears to be a distinguishing characteristic of these peculiar MSP binaries and provides a convenient way to identify additional redbacks even in the absence of radio pulsations.  The Galactic X-ray source XSS J12270--4859 (1RXS J122758.8--485343) has posed a mystery since its discovery.  Although initially classified as a cataclysmic variable (Masetti et al.~2006; Butters et al.~2008), subsequent multi-wavelength studies (Pretorius~2009; de Martino et al.~2010; Hill et al.~2011; de Martino et al.~2013; Papitto et al.~2014) revealed that this system closely resembles a quiescent LMXB.  XSS J12270--4859 is exceptional in that it was the first LMXB putatively associated with a $\\gamma$-ray source, 1FGL J1227.9--4852 \\citep{Hill11}, now known as 2FGL J1227.7--4853 \\citep{Nolan12}. Deep searches have failed to detect radio pulsation from XSS J12270--4859 during its high optical/X-ray flux state \\citep{Hill11}.  Prior to 2012 November/December, the X-ray spectrum of XSS J12270--4859 was well described by a pure powerlaw with $\\Gamma=1.7$ and a 0.1--10 keV luminosity of $\\sim$$6\\times10^{33}$ erg s$^{-1}$ \\citep{deM10,deM13}.  The source exhibited occasional intense flares and peculiar, frequent drops in X-ray flux that are not correlated with orbital phase. This behavior is reminiscent of the erratic flux variations observed in PSRs J1023+0038 (Patruno et al.~2014; Tendulkar et al.~2014; S.~Bogdanov in prep.) and J1824--2452I (Papitto et al.~2013; Linares et al.~2014) in their accretion disk states, hinting at a connection with these systems. Recent optical and X-ray observations revealed that XSS J12270--4859 had undergone a substantial decline in brightness and no longer exhibits evidence for an accretion disk (Bassa et al.~2014). The abrupt change to a disk-free system appears to have occured between 2012 November 14 and 2012 December 21. Given the similarities with PSR J1023+0038 in both states, XSS J12270--4859 is presently consistent with hosting a rotation-powered millisecond pulsar.  This prediction was confirmed with the recent detection of 1.69-millisecond radio pulsations from the system (Roy et al.~2014).  Here we present recently acquired \\textit{XMM-Newton} European Photon Imaging Camera (EPIC) and \\textit{Chandra X-ray Observatory} Advanced CCD Imaging Spectrometer (ACIS) observations of XSS J12270--4859 in its new low flux state. These observations confirm that this source has undergone a metamorphosis to a redback, i.e.~a compact binary containing a rotation-powered millisecond pulsar.  The analysis is presented as follows. In \\S 2, we provide details on the observations and data reduction procedures. In \\S3, we present an X-ray variability study. In \\S4 we summarize the phase-averaged and phase-resolved spectroscopy, while in \\S5 we report on an imaging analysis.  We provide a discussion and conclusions in \\S6. ",
        "conclusions": "As reported in Bassa et al.~(2014), the recent decline in optical and X-ray brightness and the disappearance of the previously prominent optical emission lines are indicative of the disappearance of the accretion disk in XSS J12270--4859. The \\textit{XMM-Newton} and \\textit{Chandra} observations of XSS J12270--4859 in its new low flux state presented herein reveal further information regarding this binary:  1.~Previously absent large-amplitude modulation at the binary orbital period is now present.  2.~The time-averaged X-ray luminosity of $\\sim$$(1-7)\\times 10^{32}$ erg s$^{-1}$ is smaller by an order of magnitude relative to the accretion disk dominated state.  3.~The predominantly non-thermal spectrum with spectral photon index of $\\Gamma=1.2$ is significantly harder than the $\\Gamma=1.7$ observed in the high flux state.  4.~The spectral index of the power-law emission appears not to vary substantially even though the flux changes by a factor of $\\sim$1.5--1.7 over an orbital period.  These X-ray properties, as well as the optical properties described in \\citet{Bassa14}, are identical to what is observed from ``redback'' radio MSPs, including PSRs J1023+0038 \\citep{Arch09,Arch10,Bog11b}, J1723--2837 \\citep{Faulk04,Craw13,Bog14}, J2215+5135 \\citep{Hes11,Gen13} in the field of the Galaxy, as well as PSR J0024--7204W in the globular cluster 47 Tuc \\citep{Camilo00,Freire03,Bog05}, PSR J1740--5340 in NGC 6397 \\citep{DAm01,Bog10}, and PSR J1824--2452H in M28 \\citep{Bog11a,Pal10}. Therefore, the results obtained herein, and the recent radio pulsar discovery by \\citet{Roy14} further demonstrate that XSS J12270--4859 is a redback, i.e. a compact binary with an active rotation-powered pulsar and no accretion disk.  The non-thermal X-rays from redback systems are generated in an intra-binary shock formed by the interaction of the pulsar wind with material from the close companion star \\citep{Arons93}. As shown in \\citet{Bog11b}, the decline in X-ray flux at $\\phi_b\\approx0.25$ seen in PSR J1023+0038 can be reproduced by a simple geometric model in which the intra-binary shock is partially occulted by the bloated secondary star. In this scenario, the depth and phase extent of the X-ray eclipse require that the shock be situated at or very near the surface of the secondary star on the side exposed to the pulsar wind.   There is no indication of extended X-ray emission associated with XSS J12270--4859 that could arise due to outflows in the LMXB phase or the interaction of the pulsar wind with the ambient medium.  An adjacent diffuse source is likely not associated with the binary and is probably a background galaxy cluster. Given that only two MSPs, PSRs B1957+21 \\citep{Stap03} and J2124--3358 \\citep{Hui06} show evidence for X-ray bow shocks, the absence of nebular emission associated with XSS J12270--4859 does not imply the lack of a pulsar wind.  The lack of knowledge of the orbital separation, the companion mass, and the pulsar spin-down luminosity ($\\dot{E}$) for XSS J12270--4859, do not allow meaningful constraints on the physics of the intrabinary shock. Nevertheless, given the nearly identical X-ray properties compared to other redbacks it is possible to place crude constraints on the energetics of the MSP in XSS J12270--4859. In particular, if we consider that the shock luminosity at X-ray maximum is $\\sim$$2\\times10^{-3}\\dot{E}$ for PSRs J1023+0038 \\citep{Arch10,Bog11b} and J1723--2837 \\citep{Bog14}, the implied spin-down of the pulsar in XSS J12270--4859 is $(0.7-5)\\times10^{35}$ erg s$^{-1}$ assuming a distance range $1.4-3.6$ kpc. This suggests that the pulsar in XSS J12270--4859 may be among the small number of energetic MSPs with $\\dot{E}\\gtrsim10^{35}$ erg s$^{-1}$.  The 0.1--100 GeV $\\gamma$-ray luminosity of 2FGL J1227.7--4853, the putative $\\gamma$-ray counterpart of XSS J12270--4859, of $L_{\\gamma}=(0.8-5)\\times10^{34}$ erg s$^{-1}$ provides an independent hard lower limit on the $\\dot{E}$ of the pulsar, if we assume a $\\gamma$-ray conversion efficiency of 100\\%. For efficiencies typical of the MSP population \\citep[$\\lesssim$50\\%; see Table 10 in][]{Abdo13}), this yields $\\dot{E}\\gtrsim10^{35}$ erg s$^{-1}$, in general agreement with the limit obtained from the X-ray luminosity. Continued timing of the radio pulsar would allow a determination of the true $\\dot{E}$ of this pulsar.   The redback nature of the XSS J12270--4859 binary determined from the X-ray analysis presented above implies that the pulsar should undergo radio eclipses of frequency-dependent duration around superior conjunction. These eclipses are caused by intra-binary plasma, presumably emanating from the companion star.  For most redback and black widow systems that are detectable in radio pulsations, the radio emission is eclipsed for $\\lesssim$50\\% of the orbit, although there is indication that for some systems this fraction can be close to 100\\% \\citep[see][for the case of PSR J1311--3430]{Ray13}.  The non-detections with the Parkes telescope \\citep{Bassa14} and only a single, brief detection with the GMRT \\citep{Roy14} despite deep pulsation searches, are indicative of a high radio eclipse fraction for XSS J12270--4859 that likely varies from orbit to orbit.  The large $\\dot{E}$ estimated above is consistent with this scenario, since an energetic wind would drive off material from the surface of companion at a higher rate than an MSP with a lower $\\dot{E}$. This may result in occasional severe enshrouding by this stripped material, causing the pulsar to be eclipsed at radio frequencies for a larger fraction of the orbit. This would make XSS J12270--4859 a so-called ``hidden'' MSP, as postulated by \\citet{Tav91}.  The X-ray properties of XSS J12270--4859 in both the present disk-free radio pulsar state and the accretion disk-dominated state \\citep[see][]{deM10,deM13} are virtually identical to those of PSRs J1023+0038 \\citep{Arch10,Bog11b,Pat14} and J1824--2452I \\citep{Pap13,Lin14}. This remarkable consistency in X-ray characteristics provides a powerful discriminant for identifying more such MSP binaries, especially ones that are heavily enshrouded and hence cannot be identified via radio pulsation searches."
    },
    "1402/1402.1168_arXiv.txt": {
        "abstract": "We present a statistical study of velocities of Ly$\\alpha$, interstellar (IS) absorption, and nebular lines and gas covering fraction for Ly$\\alpha$ emitters (LAEs) at $z\\simeq 2$. We make a sample of 22 LAEs with a large Ly$\\alpha$ equivalent width (EW) of $\\gtrsim 50$\\,\\AA\\ based on our deep Keck/LRIS observations, in conjunction with spectroscopic data from the Subaru/FMOS  program and the literature. We estimate the average velocity offset of Ly$\\alpha$ from a systemic redshift determined with nebular lines to be $\\Delta v_{\\rm Ly\\alpha}=234\\pm 9$ km s$^{-1}$. Using a Kolmogorv-Smirnov test, we confirm the previous claim of Hashimoto et al. (2013) that the average $\\Delta v_{\\rm Ly\\alpha}$ of LAEs is smaller than that of LBGs. Our LRIS data successfully identify blue-shifted multiple IS absorption lines in the UV continua of four LAEs on an individual basis. The average velocity offset of IS absorption lines from a systemic redshift is $\\Delta v_{\\rm IS}=204 \\pm 27$ km s$^{-1}$, indicating LAE's gas outflow with a velocity comparable to typical LBGs. Thus, the ratio, $ R^{\\rm Ly\\alpha}_{\\rm IS} \\equiv \\Delta v_{\\rm Ly\\alpha}/\\Delta v_{\\rm IS}$ of LAEs, is around unity, suggestive of low impacts on Ly$\\alpha$ transmission by resonant scattering of neutral hydrogen in the IS medium. We find an anti-correlation between Ly$\\alpha$ EW and the covering fraction, $f_c$, estimated from the depth of absorption lines, where $f_c$ is an indicator of average neutral hydrogen column density, $N_{\\rm HI}$. The results of our study support the idea that $N_{\\rm HI}$ is a key quantity determining Ly$\\alpha$ emissivity. ",
        "introduction": "Ly$\\alpha$ Emitters (LAEs) are an important population of high-$z$ star-forming galaxies in the context of galaxy formation.  LAEs at $z=2-7$ and beyond $z=7$ are found by narrow-band (NB) imaging observations based on an NB excess resulting from their prominent Ly$\\alpha$ emission \\citep[e.g., ][]{2010ApJ...711..928C,2007ApJ...667...79G,2012ApJ...744..110C,2008ApJS..176..301O,2008ApJ...677...12O,2010ApJ...723..869O,2010ApJ...725..394H,2007ApJ...660.1023F,2011ApJ...734..119K,2006ApJ...648....7K,2012ApJ...752..114S}. Observational studies on a morphology and spectral energy distribution (SED) of LAEs reveal that such a galaxy is typically young, compact, less-massive, less-dusty than other high-$z$ galaxy populations, and a possible progenitor of Milky Way mass galaxies \\citep[e.g., ][]{2011ApJ...743....9G,2011ApJ...733..114G,2010MNRAS.402.1580O,2007ApJ...671..278G,2011ApJ...740...71D,2008ApJ...681..856R,2012MNRAS.424.1672D}. Additionally, LAEs are used to measure the neutral hydrogen fraction at the reionizing epoch, because Ly$\\alpha$ photons are absorbed by intergalactic medium (IGM).  Ly$\\alpha$ emitting mechanism is not fully understood due to the highly-complex radiative transfer of Ly$\\alpha$ in the interstellar medium (ISM). Many theoretical models have predicted that the neutral gas and/or dust distributions surrounding central ionizing sources are closely linked to the Ly$\\alpha$ emissivity \\citep[e.g., ][]{1991ApJ...370L..85N, 2008ApJ...678..655F,2013ApJ...766..124L,2009ApJ...704.1640L,2007ApJ...657L..69L,2013arXiv1302.7042D,2013arXiv1308.1405Z,2010ApJ...716..574Z,2012arXiv1209.5842Y}. Thus, resonant scattering in the neutral ISM can significantly attenuate the Ly$\\alpha$ emission.  Ly$\\alpha$ emissivity may not only depend on the spatial ISM distribution, but on the gas kinematics as well. The large-scale galactic outflows driven by starbursts or active galactic nuclei could allow Ly$\\alpha$ photons to emerge at wavelengths where the Gunn-Peterson opacity is reduced, and consequently enhance the Ly$\\alpha$ emissivity, particularly in the high-$z$ Universe \\citep[e.g., ][]{2010MNRAS.408..352D}. The outflow may also blow out the Ly$\\alpha$ absorbing ISM. The gas kinematics of LAEs has been evaluated from the Ly$\\alpha$ velocity offset ($\\Delta v_{{\\rm Ly\\alpha}}$) with respect to the systemic redshift ($z_{\\rm sys}$) traced by nebular emission lines (e.g, H$\\alpha$, $[$O {\\sc iii}$]$) from their H {\\sc ii} regions. In the past few years, deep NIR spectroscopic studies have detected nebular emission lines from $\\sim10$ LAEs at $z=2-3$, and measured their $\\Delta v_{{\\rm Ly\\alpha}}$ \\citep{2011ApJ...730..136M, 2013ApJ...765...70H,2013A&A...551A..93G,2011ApJ...729..140F,2013ApJ...775...99C}. The Ly$\\alpha$ emission lines for these LAEs are redshifted from their $z_{\\rm sys}$ by a $\\Delta v_{{\\rm Ly\\alpha}}$ of $200-300$ km s$^{-1}$. \\citet{2013ApJ...765...70H} find an anti-correlation between Ly$\\alpha$ equivalent width (EW) and $\\Delta v_{{\\rm Ly\\alpha}}$ in a compilation of LAE and LBG samples. This result is in contrast to a simple picture where Ly$\\alpha$ photons more easily escape in the presence of a galactic outflow.  However, the Ly$\\alpha$ velocity offset is thought to increase with both resonant scattering in H {\\sc i} gas clouds as well as galactic outflow velocity \\citep[e.g., ][]{2006A&A...460..397V,2008A&A...491...89V}. The anti-correlation could result from a difference in H {\\sc i} column density ($N_{\\rm HI}$) rather than outflowing velocity. The gas kinematics can be investigated more directly from the velocity offset between interstellar (IS) absorption lines of the rest-frame UV continuum and $z_{\\rm sys}$ (IS velocity offset; $\\Delta v_{{\\rm IS}}$). The IS velocity offset traces the speed of outflowing gas clouds, and may help to distinguish the two effects on $\\Delta v_{{\\rm Ly\\alpha}}$.  For UV-continuum selected galaxies, the $\\Delta v_{{\\rm IS}}$ has been measured for $>100$ objects \\citep[e.g.][]{2001ApJ...554..981P,2012MNRAS.427.1973C,2012ApJ...745...33K,2013ApJ...777...67S,2010ApJ...717..289S}. \\citet{2010ApJ...717..289S} find that LBGs have an average of $\\langle\\Delta v_{{\\rm IS}}\\rangle = -164$ km s$^{-1}$ in their sample of 89 LBGs at $z\\sim3$. This statistical study indicates ubiquitousness of galactic outflow in LBGs. However, there have been no NB-selected galaxies with a $\\Delta v_{{\\rm IS}}$ measurement to date except for a stacked UV spectrum in \\citet{2013ApJ...765...70H}. This is because it is difficult to estimate $\\Delta v_{{\\rm IS}}$ for individual LAEs, especially for galaxies with a large Ly$\\alpha$ EW of $\\gtrsim 50$\\,\\AA\\, due to their faint UV-continuum emission, while $\\Delta v_{{\\rm IS}}$ are measured for some UV-selected galaxies with EW(Ly$\\alpha) \\sim50$\\,\\AA\\, \\citep[e.g., ][]{2010ApJ...719.1168E}. A statistical investigation of Ly$\\alpha$ kinematics for LAEs could shed light on the physical origin of the anti-correlation and the underlying Ly$\\alpha$ emitting mechanism.    This is the second paper in the series exploring the Ly$\\alpha$ emitting mechanisms\\footnote{The first paper presents a study on LAE structures \\citep{2014ApJ...785...64S}.}. In this paper, we present the results of our optical and NIR spectroscopy for a large sample of $z=2.2$ LAEs with Keck/LRIS and Subaru/FMOS to verify possible differences of $\\Delta v_{{\\rm Ly\\alpha}}$ and $\\Delta v_{{\\rm IS}}$ between LAEs and LBGs. These spectroscopic observations are in an extension of the project of \\citet{2013ApJ...765...70H} aiming to confirm the anti-correlation between Ly$\\alpha$ EW and $\\Delta v_{{\\rm Ly\\alpha}}$. The organization of this paper is as follows. In Section \\ref{sec_targets}, we describe the details of the LAEs targeted for our spectroscopy. Next, we show our optical and NIR spectroscopic observations in Section \\ref{sec_observation}. We present methods to reduce the spectra, and to measure kinematic quantities such as $\\Delta v_{{\\rm Ly\\alpha}}$ and $\\Delta v_{{\\rm IS}}$ in \\ref{sec_reduction}. We perform SED fitting to derive physical properties in Section \\ref{sec_sed_fitting}. We compare kinematic properties between LAEs and LBGs in Section \\ref{sec_results}, and discuss physical origins of possible differences in these quantities in Section \\ref{sec_discussion}. In the last section Section \\ref{sec_conclusion}, we summarize our findings.  Throughout this paper, we adopt the concordance cosmology with $(\\Omega_m, \\Omega_\\Lambda, h)=(0.3, 0.7, 0.7)$, \\citep{2011ApJS..192...18K}. All magnitudes are given in the AB system \\citep{1983ApJ...266..713O}. ",
        "conclusions": "We carry out deep optical spectroscopy for our large sample of LAEs at $z=2.2$ in order to detect their Ly$\\alpha$ lines with Keck/LRIS. We compare redshifts of the Ly$\\alpha$ and nebular emission lines detected with Subaru/FMOS, and calculate $\\Delta v_{\\rm Ly\\alpha}$ for new 11 LAEs. This observation doubles the sample size of LAEs with a $\\Delta v_{\\rm Ly\\alpha}$ measurement in literatures.  The conclusions of this study are summarized below.  \\begin{itemize} \\item Almost all of our new LAEs have a $\\Delta v_{\\rm Ly\\alpha}$ of $\\sim200$ km s$^{-1}$ which is systematically-smaller than that of LBGs. Using 22 LAEs with $\\Delta v_{\\rm Ly\\alpha}$ measurements taken from our new observations and the literature, we definitively confirm the anti-correlation between Ly$\\alpha$ EW and $\\Delta v_{\\rm Ly\\alpha}$ suggested by previous work. \\item Long exposure times and the high sensitivity of LRIS at blue wavelengths enabled us to successfully detect IS absorption lines against faint UV continua from four individual LAEs. These IS absorption lines are found to be blueshifted from the systemic redshift by $200-300$ km s$^{-1}$, indicating strong gaseous outflows are present even in LAEs. \\item We estimate $R_{\\rm IS}^{\\rm Ly\\alpha}$ ($\\equiv\\Delta v_{\\rm Ly\\alpha}/\\Delta v_{\\rm IS}$) that would be a quantity sensitive to $N_{\\rm HI}$ for the four UV continuum-detected LAEs. We find the value of $R_{\\rm IS}^{\\rm Ly\\alpha}$ in LAEs to be smaller than that of LBGs, indicating a lower $N_{\\rm HI}$ in LAEs. We performed a test for correlations between $R_{\\rm IS}^{\\rm Ly\\alpha}$ and physical properties inferred from SED fitting. As a result, we tentatively conclude that SFR may be most closely related to $R_{\\rm IS}^{\\rm Ly\\alpha}$. The correlation may suggest that the star formation preferentially occurs in systems with large amounts of neutral hydrogen gas, which would have a larger value of $R_{\\rm IS}^{\\rm Ly\\alpha}$. \\item We estimate the covering fraction, $f_c$, of surrounding H {\\sc i} gas from the depth of LIS absorption lines the four LAEs. We identify a tentative trend for $f_c$ to decrease with increasing Ly$\\alpha$ EW, as suggested by a study for LBGs in \\citet{2013ApJ...779...52J}. A central source being covered by patchy H {\\sc i} gas clouds would lead to a small flux-averaged $N_{\\rm HI}$ corresponding to a small $R^{\\rm Ly\\alpha}_{\\rm IS}$. In this condition, Ly$\\alpha$ photons could easily escape less affected by resonant scattering in the clouds. \\item The results of our kinematic analyses support the idea that the H {\\sc i} column density is a key quantity determining Ly$\\alpha$ emissivity.  \\end{itemize}  In this kinematic study, we obtain $\\Delta v_{\\rm IS}$, $R_{\\rm IS}^{\\rm Ly\\alpha}$, and $f_c$ only for objects with a moderate Ly$\\alpha$ EW of $20-100$\\,\\AA\\, which overlaps with the Ly$\\alpha$ EW range of LBG samples in e.g., \\citet{2003ApJ...588...65S}. We need to estimate these quantities for objects with a higher Ly$\\alpha$ EW in order to check whether such objects follow the kinematic trends found in this study."
    },
    "1402/1402.4051_arXiv.txt": {
        "abstract": "{Super-Alfv\\'enic shocks associated with coronal mass ejections (CMEs) can produce radio emission known as Type II bursts. In the absence of direct imaging, accurate estimates of coronal electron densities, magnetic field strengths, and Alfv\\'en speeds are required to calculate the kinematics of shocks. To date, 1D radial models have been used, but these are not appropriate for shocks propagating in non--radial directions.} {Here, we study a coronal shock wave associated with a CME and Type II radio burst using 2D electron density and Alfv\\'en speed maps to determine the locations that shocks are excited as the CME expands through the corona.}{Coronal density maps were obtained from emission measures derived from the Atmospheric Imaging Assembly (AIA) on board the \\emph{Solar Dynamic Observatory} ($SDO$) and polarized brightness measurements from the Large Angle and Spectrometric Coronagraph (LASCO) on board the \\emph{Solar and Heliospheric Observatory} ($SOHO$). Alfv\\'en speed maps were calculated using these density maps and magnetic field extrapolations from the Helioseismic and Magnetic Imager ($SDO$/HMI). The computed density and Alfv\\'en speed maps were then used to calculate the shock kinematics in non-radial directions.}{Using the kinematics of the Type II burst and associated shock, we find our observations to be consistent with the formation of a shock located at the CME flanks where the Alfv\\'en speed has a local minimum.}{The 1D density models are not appropriate for shocks that propagate non--radially along the flanks of a CME. Rather, the 2D density, magnetic field and Alfv\\'en speed maps described here give a more accurate method for determining the fundamental properties of shocks and their relation to CMEs.} ",
        "introduction": "Solar flares and coronal mass ejections (CMEs) are energetic manifestations of the restructuring coronal magnetic fields. As a CME travels through the corona, its velocity can become sufficiently larger than the background coronal Alfv\\'en speed, causing a shock wave to form along its leading edge and/or flanks \\citep{Cho2007}. It is within these shocks that electrons can be accelerated to near-relativistic energies to produce Type II radio signatures in low-frequency dynamic spectra \\citep[see,][]{Carley2013}. Although CMEs and Type II radio bursts have been studied for many decades \\citep[see,][]{Pick2006}, there remains unanswered questions in relation to where CME shocks are formed and how these phenomena are associated with the generation of Type II bursts.  It has long been suggested that Type II radio bursts are signatures of coronal shocks \\citep{Wild1950,Uchida1960}. They appear as features slowly drifting toward lower frequencies at decimetric to kilometric wavelengths in dynamic radio spectra. This drift is the result of plasma emission generated by a super-Alfv\\'enic shock traveling upwards in the corona \\citep{Cane1981}, where density decreases with height. When direct low-frequency imaging is not available, a key problem is therefore the accurate calculation of the coronal density and Alfv\\'en speed distributions with height; This is to relate the Type II emission frequency to its height and to investigate the direction of the shock propagation. \\begin{figure}[h] \\begin{center} \\includegraphics[trim=1cm 1cm 1cm 1cm,clip=true,width=9cm,angle=0]{plot_density-eps-converted-to.pdf} \\caption{A variety of commonly used electron density models, which can give very different height estimates. For example, a density of $10^7$ cm$^{-3}$ can be located at $\\sim$1.4~$R_{\\odot}$ or $\\sim$2.0~$R_{\\odot}$, depending on the density model. Plotted for comparison, the set of density measurements obtained with Mars Express (MEX) between the years 2006--2008 \\citep{Verma2013}.} \\label{Fig_1_models} \\end{center} \\end{figure} The plasma frequency is related to the density of the emitting plasma by $f_\\mathrm{p}= C\\sqrt{n_\\mathrm{e}} $, where $C=8980$ Hz cm$^{3/2}$ is a constant. To derive the shock kinematics, electron density models are normally employed to relate the plasma density to its coronal height and velocity. Specifically, the shock radial velocity is related to the plasma frequency drift rate, $\\mathrm{d} f_\\mathrm{p}/\\mathrm{d} t$,  and the electron density model, $ n_\\mathrm{e}(r)$,  by, \\begin{equation} v= \\frac{2\\sqrt{n_\\mathrm{e}}}{C} \\left( \\frac{\\mathrm{d} n_\\mathrm{e}}{\\mathrm{d} r} \\right)^{-1} \\frac{\\mathrm{d} f_\\mathrm{p}}{\\mathrm{d} t} , \\end{equation} \\begin{figure} \\centering \\includegraphics[trim=0cm 0cm 0cm 0cm, clip=true, width=9cm, angle=0]{Fig_2_one_column-eps-converted-to.pdf} \\caption{\\emph{GOES--15} soft X--ray light curve showing an X1.4 flare starting at 10:29:00~UT on 2011 Septemeber 22 (top) and the associated RSTO dynamic spectrum showing a Type II radio burst starting at 10:39:06~UT (bottom). This burst shows fundamental (F) and harmonic (H) emission.} \\label{RSTO} \\end{figure} where $v$ is the shock velocity, $n_\\mathrm{e}$ is the coronal plasma electron density and $r$ is the heliocentric radial distance. Different coronal density models can therefore lead to differing kinematics. Several density models have been used to calculate the Type II shock position, such as the \\citet{Newkirk1961} model derived from the barometric height behavior of a gravitationally stratified corona, the \\citet{Saito1977} model obtained from measurements of coronal polarized brightness ($pB$), and the \\citet{Mann1999} model, derived from solutions of magnetohydrodynamic equations.  The use of an arbitrary radial density model can lead to an inaccurate calculation of shock heights and hence velocities. This is due to the significant difference in the shock height derived from different density models. An example is shown in Fig.~\\ref{Fig_1_models}, where an electron density of $10^7$~cm$^{-3}$, occurs at a height of 1.4~$R_{\\odot}$ for the \\citet{Mann1999} model,  while occurs at a height of 2.0~$R_{\\odot}$ for the Allen-Baumbach model \\citep{Allen1947}. Another reason for an inaccurate shock height and velocity calculation is the time variability of the coronal density distribution \\citep{Parenti2000,Bemporad2003}, which is not taken into account with typically employed density models. Finally, if a Type II radio source propagates non-radially, speeds derived from 1D radial density models underestimate the true shock speed.  Knowledge of the coronal density in the 2D plane is therefore crucial for determining accurate shock kinematics, while the knowledge of the 2D Alfv\\'en speed is important to determine when the propagating shock reaches a super-Alfv\\'enic speed. A 2D analytic model of the Alfv\\'en speed was presented by \\citet{Warmuth2005}. However, due to the temporal variability of the density and magnetic field in the corona, it is important to use observational data specific to the radio emission time rather than a generic analytic model for the Alfv\\'en speed. \\begin{figure*} \\begin{center} \\includegraphics[trim=6.1cm 2.5cm 6.5cm 2.9cm,clip=true,width=6.7cm,angle=90]{sdo_running_diff_cme_new-eps-converted-to.pdf} \\caption{{\\it SDO}/AIA running--difference images of the erupting CME at 10:36:39--10:40:24~UT in the 211~\\AA\\ passband. The CME leading edge is outside the field--of--view of $SDO$ in panel (a), while the propagation of the coronal bright front is visible at the location of the CME flanks at 10:38:38~UT panel (b) and at 10:40:24~UT panel (c).} \\label{SDO_cme} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[trim=6cm 3.8cm 5.9cm 4cm,clip=true,width=7.9cm,angle=90]{Intensity_Density_SDO_new-eps-converted-to.pdf} \\caption{{\\it SDO}/AIA 211~\\AA\\ passband intensity at 10:39:24~UT (a) and the corresponding density map (b). The contours show the harmonic plasma emission frequencies related to the measured coronal densities at 150, 200 and 300~MHz. The contours have been smoothed by a boxcar average of 3.6 arcsecs for display purposes, while} dashed lines are at 1.0~$R_{\\odot}$ and 1.2~$R_{\\odot}$. \\label{Density_map_sdo} \\end{center} \\end{figure*}  Here, a new method to calculate coronal densities, magnetic field strengths, and Alfv\\'en speeds in a 2D plane is presented. These 2D maps are then used to calculate the kinematics and the direction of propagation of a Type II radio burst observed on 2011 September 22. The Type II radio burst and coronal mass ejection (CME) are described in Section~\\ref{observations}.  The observational method and the models used to produce the 2D density maps are described in Section~\\ref{density}, while the Alfv\\'en maps are shown in Section~\\ref{Alfven}. Results are presented in Section~\\ref{results}, and finally, conclusions are discussed in Section~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}   A new method to obtain semi--empirical 2D maps of coronal density and Alfv\\'en speed has been presented. Using these calculated maps, a decametric Type II burst and a CME that occurs on 2011 September 22, were analyzed to understand their relationship.  The analysis of the Type II kinematics, using the 2D density maps, have been performed for non--radial profiles in the POS starting from the active region. Previously, analytical radial density models have been employed to calculate the radio burst kinematics \\citep[e.g., ][]{Robinson1985,Vrsnak2002}. Using a 2D density map at the specific time of the radio burst  and performing a non--radial analysis, we were able to relate the Type II burst with the CME and discriminate which portion of the CME front was responsible for triggering the Type II emission. Comparing the CME kinematics calculated for different non-radial traces to the Type II burst we find evidence for the shock to be formed in the flank region of the CME. This agrees with previous interpretations \\citep[e.g.,][]{Classen2002,Mancuso2004}. By calculating the height of the upstream shock front (70~MHz) using the 2D density map and relating it to the position of the CME, we were able to locate the position of the shock in the CME flanks at a height of $\\sim$1.6~$R_{\\odot}$.  In addition, we were able to compare the Type II burst shock speed with the local Alfv\\'en speed using the 2D Alfv\\'en speed maps . The Type II burst speed was found to be super-Alfv\\'enic \\citep[e.g.,][]{Vrsnak2008,Cho2005,Klein1999}). Furthermore, the decelerating phase of the Type II shock can be related with the fading of the Type II burst in the dynamic spectrum as the shock approaches the local Alfv\\'en speed.  The method used to construct 2D density and Alfv\\'en speed maps presented in this work can be used for the study of all Type II radio bursts. These time specific semi-empirical maps improve the calculation of shock kinematics and represent a step forward from the standardized radial density profiles currently employed. Further investigation relating the kinematics of the limb event of Type IIs and CMEs using the 2D maps and radio imaging is necessary. This can be done with Type II radio bursts in the NRH frequency range (i.e., 150--445~MHz) where direct radio imaging can be used to locate the position of the burst, or with lower frequency imaging telescopes, such as the LOw Frequency ARray \\citep[LOFAR;][]{VanHaarlem2013}."
    },
    "1402/1402.4267_arXiv.txt": {
        "abstract": "The dispersion of small amplitude, impulsively excited wave trains propagating along a magnetic flux tube is investigated. The initial disturbance is a localized transverse displacement of the tube that excites a fast kink wave packet. The spatial and temporal evolution of the perturbed variables (density, plasma displacement, velocity, \\ldots) is given by an analytical expression containing an integral that is computed numerically. We find that the dispersion of fast kink wave trains is more important for shorter initial disturbances (i.e. more concentrated in the longitudinal direction) and for larger density ratios (i.e. for larger contrasts of the tube density with respect to the environment density). This type of excitation generates a wave train whose signature at a fixed position along a coronal loop is a short event (duration $\\simeq 20$~s) in which the velocity and density oscillate very rapidly with typical periods of the order of a few seconds. The oscillatory period is not constant but gradually declines during the course of this event. Peak values of the velocity are of the order of 10~km~s$^{-1}$ and are accompanied by maximum density variations of the order of 10--15\\% the unperturbed loop density. ",
        "introduction": "In the last two decades abundant evidence about waves and oscillations in the solar atmosphere has been gathered. Events of various nature (standing and propagating waves) and in various environments (chromosphere, prominences, active regions, coronal holes) have been detected. Here we present some examples, while emphasizing that the following list is not exhaustive. In the chromosphere, propagating and standing transverse waves have been detected in spicules \\citep{zaqarashvili2007,okamoto2011}, in mottles \\citep{kuridze2012}, and in active region fibrils \\citep{pietarila2011}. Regarding solar prominences, \\citet{lin2007} used high-resolution H$\\alpha$ filtergrams and observed traveling transverse waves in thin filament threads. \\citet{okamoto2007} found transverse oscillations of flowing active region filament threads observed with {\\em Hinode} SOT. Oscillatory events of different nature have been observed in coronal loops: transverse oscillations of active region loops triggered by a disturbance that propagates from the central flare site \\citep{aschwanden1999,nakariakov1999}; high-frequency, compressible waves traveling along an active region coronal loop \\citep{williams2001,williams2002}; Doppler shift oscillations caused by waves propagating in the upper part of coronal loops \\citep{tian2012}; ubiquitous waves in the solar corona propagating upwards along magnetic field lines \\citep{tomczyk2007,tomczyk2009}; etc. Reviews about waves and oscillations in spicules, prominences, and coronal structures can be found in \\citet{zaqarashvili_erdelyi2009}, \\citet{arregui2012}, \\citet{nakariakov_verwichte2005}, and \\citet{demoortel_nakariakov2012}.  Many of these events are examples of magnetic flux tubes being perturbed by an external agent. The transverse loop oscillations described by \\citet{aschwanden1999} and \\citet{nakariakov1999} are a remarkable case because time series of EUV images allow to see the lateral swaying of a coronal loop. This particular phenomenon has been interpreted as a standing fast kink mode oscillation. Fast kink (i.e. transverse) waves propagating along magnetic flux tubes are more frequent in the literature than their standing oscillation counterparts. They have been observed not only in coronal loops \\citep[e.g.][]{tian2012}, but also in spicules \\citep{zaqarashvili2007}, in mottles \\citep{kuridze2012}, in active region fibrils \\citep{pietarila2011}, and in filament threads \\citep{lin2007,okamoto2007}. What is common to all these phenomena is that an external excitation causes the transverse displacement of a magnetic tube and that this perturbation propagates along the tube, where it is usually detected as a time variation of the Doppler velocity or the magnetic tube position. These propagating waves can be excited, e.g., by a periodic driver acting at a fixed position of the magnetic tube. This has been the mechanism invoked by \\citet{VTG2010} and \\citet{TGV2010} to explain the waves propagating along the coronal magnetic field observed by \\citet{tomczyk2007} and \\citet{tomczyk2009}. A fast kink wave propagating along a uniform, cylindrical magnetic tube produces a periodic transverse motion of the tube such as that in Figure~\\ref{fig_cylinder_kink}(a).   \\begin{figure}[ht!] \\centerline{(a)\\includegraphics[width=0.3\\textwidth,angle=0]{f1a.eps}\\hspace{1cm}(b)\\includegraphics[width=0.3\\textwidth,angle=0]{f1b.eps}} \\caption{Uniform cylindrical magnetic flux tube subject to (a) a fast kink mode (i.e. transverse, periodic along the tube) perturbation and (b) a localized transverse perturbation. (A color version of this figure is available in the online journal.)} \\label{fig_cylinder_kink} \\end{figure}   Propagating transverse waves can also be generated by an impulsive, concentrated transverse displacement of the magnetic tube (see Figure~\\ref{fig_cylinder_kink}(b)). Such a perturbation of the magnetic tube is a superposition of fast kink eigenmodes, each with its own amplitude. In the absence of dispersion, the initial wave form would keep its original shape during its propagation, that is, the bulge of Figure~\\ref{fig_cylinder_kink}(b) would simply propagate unaltered along the cylinder. Fast kink eigenmodes of a magnetic tube are dispersive, however, and so as time evolves the initial hump transforms into an oscillatory train containing several maxima and minima. In other words, away from the excitation point the magnetic tube can suffer not one, but several lateral oscillations about its equilibrium position as the wave train passes by. Hence, transverse oscillations of magnetic tubes do not necessarily require the presence of a continuous driver.  The purpose of this work is to study the dispersion of linear fast kink wave trains propagating along uniform, cylindrical magnetic flux tubes. These wave trains are excited by an impulsive, localized initial disturbance (Figure~\\ref{fig_cylinder_kink}(b)). To solve this initial value problem we use a technique based on the method of Fourier integrals \\citep[cf.][Section~11.2]{whitham1974}, that consists of expressing the initial perturbation as a sum of eigenmodes. This means that the properties of the magnetic tube eigenmodes need to be well known in advance, and for this reason their main features are summarized in Sections~\\ref{sect_normal_modes} and \\ref{sect_cont_modes}. Next the method of Fourier integrals is applied to the cylindrical magnetic tube, and an analytical expression describing the spatial and temporal variation of  perturbed variables is obtained (Section~\\ref{sect_fourier}). This formula is expressed in terms of an integral that contains contributions from all eigenmodes, with amplitudes that depend exclusively on the initial conditions. Accurate approximations to this integral are computed numerically. The propagation of a concentrated transverse disturbance is considered in Section~\\ref{subsect_numerics}, paying special attention to the dispersion of the wave train during its propagation. An application to coronal loop transverse oscillations is carried out in Section~\\ref{sect_loop_application} and a discussion of the results and our conclusions are presented in Section~\\ref{sect_conclusions}. ",
        "conclusions": "\\label{sect_conclusions}  We have analyzed the propagation of impulsively excited, linear fast kink wave packets in cylindrical magnetic tubes. The method we use is based in decomposing the initial perturbation in a sum of eigenmodes. We have imposed a concentrated transverse disturbance of the magnetic tube and have found that most of its energy is imparted to the global kink mode. The concentrated transverse impulse is the sum of many different global fast kink modes with their own amplitude and that travel at their own group velocity in the packet. For this reason, the wave train dispersion is determined by two key ingredients: on one hand, the dependence of the group velocity on the longitudinal wavenumber, that is represented in Figures~\\ref{fig_dr_over}(c) and \\ref{fig_cg_vae}. On the other hand, the amplitude of the global kink mode as a function of the longitudinal wavenumber (Figure~\\ref{fig_amplr}). We have found that large internal to external density ratios and especially shorter initial perturbations enhance wave dispersion.  The signature of a wave train at a fixed point in a coronal loop has been investigated using the radial velocity and the density perturbation as measurable quantities. We have shown that both variables display very fast oscillations with periods varying between 1 and 4~s and with detectable amplitudes. These values have been obtained for a loop radius $a=250$~km and an internal Alfv\\'en speed $\\vai=500$~km~s$^{-1}$, although these are just illustrative values. If the loop radius is doubled or its internal Alfv\\'en speed is halved, then the obtained periods and time scales are doubled. That is, both the horizontal and vertical scales in Figures~\\ref{fig_vae2_delta1_z}(b), (d) and \\ref{fig_vae2_delta1_rho_z}(b) are doubled. In view of the diversity of both $a$ and $\\vai$ in coronal loops, we conclude that the dispersion of a transverse coronal loop perturbation can be detected as periodic variations of the velocity and density with periods ranging from a fraction of a second to tens of seconds.  \\citet{murawski1998} performed a numerical simulation of impulsively generated, small-amplitude waves in a slab coronal loop model. Regarding transverse oscillations, these authors concluded that they could be detected as changes in the loop position through difference images of time series taken with a cadence of seconds or better. Moreover, \\citet{selwa2004} solved the same problem for a cylindrical loop geometry and obtained that time scales of the fast kink wave are of the order of a dozen of seconds, in agreement with the results presented here. \\citet{selwa2004} did not describe the dispersive properties of localized transverse loop disturbances and, for a perturbation centered in the loop axis, did not obtain large density variations. This contradicts our conclusions (see Figure~\\ref{fig_vae2_delta1_rho_z}), probably because the initial disturbance of \\citet{selwa2004} has a very small amplitude, of the order of 0.1~km. Impulsively generated transverse waves in coronal loops modeled as straight or curved slabs \\citep{selwa2006,selwa2007,rial2013} or as curved cylinders \\citep{selwa2011a,selwa2011b} have also been investigated, although the main aim of these studies is the analysis of transverse oscillations in line-tied loops, rather than the propagation of localized disturbances.  A localized axisymmetric compression (rather than the localized transverse excitation investigated here) is the sum of fast sausage eigenmodes, which in a uniform cylinder are highly dispersive. Studies of propagating fast sausage wave trains have been performed for example by \\citep{roberts1984}. These authors showed that an impulsive axisymmetric excitation of a magnetic cylinder generates a clear time signature when the wave packet arrives at a point along the tube away from the initial disturbance. This time signature consists of a periodic phase, followed by a quasi-periodic phase, and a decaying phase \\citep[see also][]{murawski1993b,murawski1993d,nakariakov2004,nakariakov2009}. This complex temporal behavior is produced by the dispersive character of the fast sausage mode, that results in different frequencies arriving at different times at the detection position. In the study of \\citet{nakariakov2004}, the detected density perturbation contains a continuous variation of the period and its wavelet diagram shows large power distributed between 3 and 10~s, with a total duration of the oscillatory event of about 20~s. In fact, the wavelet spectra of Figure~\\ref{fig_vae2_delta1_z}(d) and Figures~3 to 5 of \\citet{nakariakov2004} are remarkably similar in spite of the different nature of the propagating disturbances. Fast sausage wave trains propagating along coronal loops have been invoked by \\citet{nakariakov2004,nakariakov2005} as the cause of the compressible waves analyzed by \\citet{williams2001,williams2002} and \\citet{katsiyannis2003}. These waves have also been put forward as the cause of the ``wavelet tadpoles'' found in solar decimetric radio bursts \\citep{meszarosova2009,meszarosova2011,karlicky2013}. But in view of the similar density perturbations produced by a transverse excitation and an axisymmetric compression, these conclusions may need additional work to be confirmed or refuted.  Our method for solving Equations~(\\ref{eq_mass})--(\\ref{eq_induc}) in a magnetic cylinder has no underlying assumptions, except that the magnetic tube and the environment are uniform and that the internal to external density ratio is larger than one. It has several advantages over the direct numerical solution of Equations~(\\ref{eq_mass})--(\\ref{eq_induc}): the numerical integration of this set of partial differential equations requires using a two-dimensional mesh of points in the $r$- and $z$-directions. Moreover, to obtain the solution at a given time, $t$, the temporal evolution of the five perturbed variables $\\xi_r$, $\\xi_\\varphi$, $b_r$, $b_\\varphi$, and $P$ must be computed from $t=0$ to $t$. To make things worse, to produce the results of Figure~\\ref{fig_vae2_delta1_t}, for example, one needs to consider, at least, the large range of distances $0\\leq z\\leq 160a$, with $a$ the magnetic tube radius. The method we use is a numerical one and so it also suffers from numerical inaccuracy. It is nevertheless optimised to study wave propagation along a uniform magnetic tube. Imposing the initial conditions allows to obtain the amplitudes of the eigenmodes ($A_j^\\pm(k)$ from Equation~(\\ref{A_pm-main}) and $A_\\omega^\\pm(k)$ from Equation~(\\ref{A_pm_om-main})). Once these amplitudes are known, any perturbed variable can be computed at any position and time with Equation~(\\ref{Four_trans_inv-main}). A similar approach was used by \\citet{terradas2007}, based on the work of \\citet{ruderman2006}, to estimate the energy deposited in the normal modes given an external perturbation located in the solar corona.  The method can be extended to study the propagation of, for example, an axisymmetric compression of the magnetic tube. Such a disturbance can be expressed as the sum of fast sausage eigenmodes. This study will be almost a complete repetition of the analysis in this paper. The only difference is that, in the case of fast sausage waves, there is a cut-off wavenumber even for the mode fundamental in the radial direction. This implies that we can expect more substantial contribution from improper eigenmodes.  In this paper we have considered a magnetic tube with a sharp boundary. If, instead, we consider a tube with the density continuously varying from its value inside the tube to a lower value outside, then there will be wave damping due to resonant absorption. Even in this case the initial value problem can be solved using the same method as the one used in this paper. The only difference is that, in the case of a tube inhomogeneous in the radial direction, it is usually not possible to calculate the proper and improper eigenfunctions analytically, so they have to be calculated numerically."
    },
    "1402/1402.7358_arXiv.txt": {
        "abstract": "Effective WIMP models are minimal extensions of the standard model that explain the relic density of dark matter by the ``WIMP miracle.'' In this paper we consider the phenomenology of effective WIMPs with trilinear couplings to leptons and a new ``lepton partner'' particle. The observed relic abundance fixes the strength of the cubic coupling, so the parameters of the models are defined by the masses of the WIMP and lepton partner particles. This gives a simple parameter space where collider and direct detection experiments can be compared under well-defined physical minimality assumptions. The most sensitive collider probe is the search for leptons + MET, while the most sensitive direct detection channel is scattering from nuclei arising from loop diagrams. Collider and direct detection searches are highly complementary: colliders give the only meaningful constraint when dark matter is its own antiparticle, while direct detection is generally more sensitive if the dark matter is not its own antiparticle. ",
        "introduction": "\\label{sec:intro} The observation of dark matter is an unambiguous discovery of physics beyond the standard model. The simplest and most compelling explanation for dark matter is a stable thermal relic, a weakly interacting massive particle (WIMP). Assuming that it freezes out by annihilation to standard model particles via dimensionless order-1 couplings, the correct relic density is obtained for dark matter masses of order 100~GeV to 1~TeV. This ``WIMP miracle'' means that dark matter can potentially be directly produced and studied at high-energy collider experiments such as the LHC. At the same time, this coupling allows for direct detection of astrophysical dark matter through its interactions with ordinary matter. % It is one of the most % ambitious dreams of particle physics to be able to discover dark matter in both types of experiment, and make a direct comparison to properties of dark matter inferred from cosmology and astrophysics.  There are many well-motivated models of physics beyond the standard model that have a WIMP candidate (for example, the minimal supersymmetric standard model), but comparing collider and direct detection sensitivity is generally model-dependent due to the large number of % independent parameters in these theories. In a previous paper \\cite{Chang:2013oia}, we proposed a class of minimal models where it is assumed that the only particles beyond the standard model that are important for dark matter phenomenology consist of the WIMP and a standard model ``partner'' particle with the same gauge quantum numbers as one of the standard model particles. This allows a cubic coupling of the form \\beq\\eql{cubic} \\De \\scr{L} \\sim \\la (\\text{SM})  ( \\widetilde{\\text{SM}}) (\\text{DM}), \\eeq where DM denotes the dark matter particle and SM ($\\widetilde{\\text{SM}}$) denote the standard model field and its partner, respectively. The coupling $\\la$ is fixed by requiring the correct relic abundance, so the only parameters in this model are the masses of the WIMP and the standard model partner particle, once the spin and $CP$ properties of the dark matter is fixed. This gives a well-defined % and complete effective theory for dark matter phenomenology, motivating the name ``effective WIMP'' for this class of models. In particular, it allows direct comparison of the collider and direct detection searches for dark matter under well-defined physical minimality assumptions.    \\Ref{Chang:2013oia} analyzed the phenomenology in the case where the standard model particle and its partner are colored. In this case collider and direct detection experiments were found to be highly complementary. Similar models were also studied in \\Refs{An:2013xka, Bai:2013iqa,DiFranzo:2013vra,Papucci:2014iwa}. The major difference in our work is that we focus on the parameter space where the WIMP has the correct thermal relic abundance. This has important implications for the global picture. For example, \\Ref{Chang:2013oia} found that the coupling $\\la$ is enhanced in some regions of parameter space, leading to increased sensitivity for collider and direct detection searches. On the other hand, the relic abundance constraint eliminates regions of parameter space where indirect detection limits are important. There has also been other approaches to simplified dark matter such as $s$-channel mediators \\cite{Buchmueller:2013dya,An:2012ue,Shoemaker:2011vi} and dark matter with electroweak interactions \\cite{Cirelli:2005uq, Cheung:2013dua}.   In this paper, we consider effective WIMP models where the standard model particle is a lepton ($e, \\mu, \\tau$ or their corresponding neutrinos). In this case, the collider constraints come from the pair production of lepton partners followed by their decay into a WIMP plus a lepton or neutrino, and so the collider constraints are independent of the value of $\\la$. Once again, we find that collider and direct detection experiments are highly complementary. The direct detection constraints depend on $\\la$, with the dominant interactions coming from loop diagrams to a photon. The type of interactions allowed depend on whether the dark matter is its own antiparticle, which strongly affects the scattering rate. If the dark matter is its own antiparticle, the only meaningful constraints come from collider experiments, while in the opposite case direct detection experiments are generally more sensitive. There is a significant region of parameter space where both future collider and direct detection experiments can see a signal.  Imposing the relic abundance constraints has an important effect on the direct detection phenomenology. In particular, these constraints extend into the regime of large lepton partner masses because the coupling $\\la$ becomes large in this region to get the correct relic abundance. Imposing the relic abundance constraint also implies that constraints from indirect detection are not important.   The models that we consider are listed in Table~1. We consider couplings to left-handed lepton doublets only, since the phenomenology of right-handed leptons is essentially a subset of this case.  Note that this fixes the production rate at colliders which leads to slightly stronger constraints than the right-handed lepton partner case. We consider both the cases where the dark matter couples to only the first two generations of leptons, as well as the case where it couples to all generations. We find that the differences between them are small, as one might expect.  \\begin{table} \\begin{center} \\begin{tabularx}{.85\\textwidth}{|p{4cm}|p{4cm}|X|} \\hline \\multicolumn{2}{|c|}{Particles} & \\multirow{2}{.2\\textwidth}{\\parbox{4cm}{\\hspace{.5in} $\\mathcal{L}_{\\text{int}}$}}\\\\ \\cline{1-2} Dark matter $\\chi$ & Lepton partner $L$ &     \\\\ \\hline \\hline Majorana fermion& Complex scalar&  $\\la(\\chi \\ell)L^*+\\text{h.c.} $\\\\ \\hline Dirac fermion & Complex scalar & $\\la(\\chi \\ell)L^*+\\text{h.c.}$ \\\\ \\hline Real scalar & Dirac fermion&  $\\la (L^c \\ell)\\chi+\\text{h.c.}$  \\\\ \\hline Complex scalar & Dirac fermion & $\\la (L^c \\ell)\\chi+\\text{h.c.}$   \\\\ \\hline  \\end{tabularx} \\begin{minipage}{5.5in} \\caption{\\small Summary of the models considered in this paper. Spinors are written in 2-component notation. Here $\\ell$ is the left-handed lepton doublet of the standard model, $L$ is the lepton partner field, and $\\chi$ is the dark matter field. \\label{table:models}} \\end{minipage} \\end{center} \\end{table} ",
        "conclusions": "We have considered minimal extensions of the standard model containing electroweak singlet dark matter with renormalizable couplings to leptons and lepton ``partner'' particles. These models naturally explain the observed relic abundance of dark matter by the ``WIMP miracle.'' The strength of the coupling is fixed by requiring the correct relic abundance, so these models are parameterized only by the masses of the dark matter and the lepton partner particles. This gives a simple model where direct detection and dark matter searches at colliders can be compared under well-defined physical minimality assumptions.  Previous studies focused on effective WIMPs coupled to colored standard model particles, while this paper focuses on effective WIMPs coupled to leptons. As in the previous case, we find that collider and direct detection experiments are remarkably complementary. The collider limits depend only on the spin and masses of the particles, and not whether dark matter is its own antiparticle. The dominant constraint comes from the dilepton plus MET channel. We find that the mono-lepton channel is less sensitive in these models.  Direct detection experiments place interesting bounds only in the case where the dark matter is not its own anti-particle, {\\it i.e.}~Dirac or complex scalar dark matter. In these cases, current direct detection limits are generally stronger than collider limits, although there is an interesting region near $m_\\chi \\sim m_L/2$, where collider limits are competitive or stronger. In this region, it is possible to discover the dark matter in both future collider and direct detection experiments.  In particular, the next generation of direct detection experiments (e.g. XENON1T) will be sensitive to the bulk of parameter space.  However, if the dark matter is its own antiparticle ({\\it i.e.}~Majorana or real scalar dark matter) the direct detection bounds are well below the reach of current and next generation experiments, while the collider bounds are unaffected. In this case, leptophilic effective WIMPs can be discovered only in collider experiments.  {\\it Note: As this work was being completed \\Ref{Bai:2014osa} appeared, which analyzes the same models without the relic abundance constraint and makes  projections for lepton partner reaches at the 14~TeV LHC.}"
    },
    "1402/1402.5993_arXiv.txt": {
        "abstract": "{We study a dark energy model with non-zero anisotropic stress, either linked to the dark energy density or to the dark matter density. We compute approximate solutions that allow to characterise the behaviour of the dark energy model and to assess the stability of the perturbations. We also determine the current limits on such an anisotropic stress from the cosmic microwave background data by the Planck satellite, and derive the corresponding constraints on the modified growth parameters like the growth index, the effective Newton's constant and the gravitational slip.}    \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0322_arXiv.txt": {
        "abstract": "{The two objects \\object{1SWASP~J150822.80$-$054236.9} and \\object{1SWASP~J160156.04+202821.6} were initially detected from their SuperWASP archived light curves as candidate eclipsing binaries with periods close to the short-period cut-off of the orbital period distribution of main sequence binaries, at $\\sim$0.2~d.  Here, using INT spectroscopic data, we confirm them as double-lined spectroscopic and eclipsing binaries, in contact configuration.  Following modelling of their visual light curves and radial velocity curves, we determine their component and system parameters to precisions between $\\sim$2 and 11\\%.  The former system contains 1.07 and 0.55~$M_{\\sun}$ components, with radii of 0.90 and 0.68~$R_{\\sun}$ respectively; its primary exhibits pulsations with period 1/6 the orbital period of the system.  The latter contains 0.86 and 0.57~$M_{\\sun}$ components, with radii of 0.75 and 0.63~$R_{\\sun}$ respectively.} ",
        "introduction": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{star10lc.eps}} \\caption{SuperWASP light curve for J150822, folded at period of 22469.219~s, with binned mean curve overplotted.  A typical uncertainty range for a single observation is shown.  These fluxes correspond to a visual magnitude range of $\\sim$12.4--13.2.} \\label{star10lc} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{star28lc.eps}} \\caption{SuperWASP light curve for J160156, folded at period of 19572.136~s, with binned mean curve overplotted.  A typical uncertainty range for a single observation is shown.  These fluxes correspond to a visual magnitude range of $\\sim$14.1--14.8.} \\label{star28lc} \\end{figure}  The orbital period distribution of main sequence binary stars exhibits a fairly sharp lower limit at around 0.2~d \\citep{ruc92, ruc07, szymanski, paczynski}, the cause of which is the subject of ongoing research e.g. \\citet{stepien, stepien12, jiang}.  However, despite the inherent interest of this region of parameter space, relatively few eclipsing binaries (EB) have been discovered with periods near the cut-off point.  This motivated a search of the archive of the SuperWASP project (Wide Angle Search for Planets: \\citet{pollacco}) for EB candidates with apparent periods $<$20\\,000~s ($\\sim$0.2315~d), reported in \\citet{norton}.  53 candidates were found, 48 of which were new discoveries at the time.  A subsequent search of these candidate EBs for evidence of period change \\citep{lohr} corrected the periods of nine to values slightly greater than 20\\,000~s, but still $<$22600~s ($\\sim$0.2616~d).  A more rigorous search of the SuperWASP archive \\citep{lohr13} then detected 143 candidate EBs with periods $<$20\\,000~s, including 97 new discoveries since \\citet{norton}, and measured significant period changes in 74 candidates.  Here, spectroscopic data allow us to confirm two of these candidates as double-lined EBs in contact configuration (W\\,UMa-type variables): \\object{1SWASP~J150822.80$-$054236.9} (J150822) and \\object{1SWASP~J160156.04+202821.6} (J160156).  Both were initially identified in \\citet{norton}; the period of J150822 was revised upwards to 22469.2~s in \\citet{lohr} and so it did not appear in \\citet{lohr13}.  We report system and component parameters obtained for these EBs by simultaneous modelling of their SuperWASP light curves and radial velocities.  These should be of interest for the study of low-mass dwarfs and W\\,UMa systems in general, and of very short period binaries specifically. ",
        "conclusions": "J150822 and J160156 are established to be spectroscopic double-lined and eclipsing binary systems in contact configuration, composed of low-mass dwarfs.  J150822 has been modelled as consisting of 1.07 and 0.55~$M_{\\sun}$ components (mass ratio 0.51), and J160156 as having 0.86 and 0.57~$M_{\\sun}$ components (mass ratio 0.67).  The primary of J150822 appears to be pulsating with a period 1/6 of the orbital period.  Both systems are plausibly undergoing mass transfer; this may be related to the primaries' radii being smaller, and the secondaries' radii being larger, than would be typical for single stars with these masses.  The parameters obtained here should contribute to our understanding of low-mass stars and contact binary systems, since relatively few binaries are known with such short orbital periods.  We hope to follow up further candidate short-period EBs listed in \\citet{lohr13} with multi-colour photometry and spectroscopy, with a view to confirming their binary nature.  Many of them should be good prospects for full solution, and capable of significantly extending our knowledge of this aspect of the field."
    },
    "1402/1402.5507_arXiv.txt": {
        "abstract": "Eclipsing millisecond pulsars in close ($P_b < 1$~day) binary systems provide a different view of pulsar winds and shocks than do isolated pulsars. Since 2009, the numbers of these systems known in the Galactic field has increased enormously. We have been systematically studying many of these newly discovered systems at multiple wavelengths. Typically, the companion is nearly Roche-lobe filling and heated by the pulsar which drives mass loss from the companion. The pulsar wind shocks with this material just above the surface of the companion. We discuss various observational properties of this shock, including radio eclipses, orbitally modulated X-ray emission, and the potential for $\\gamma$-ray emission. Redbacks, whose companions are likely non-degenerate and significantly more massive, generally have more luminous shocks than black widows which have very low mass companions. This is expected since the more massive redback companions intercept a greater fraction of the pulsar wind. We also compare these systems to accreting millisecond pulsars, which may be progenitors of black widows and in some cases can pass back and forth between redback and accretion phases. ",
        "introduction": "In 1986, the first binary pulsar to exhibit radio eclipses, PSR B1957+20, was discovered at the Arecibo observatory (Fruchter et al. 1988a). This 1.61~ms pulsar is in a 9.2~hr orbit around a very low mass ($M_c \\sim 0.02 M_{\\odot}$) companion. Regular radio eclipses were observed for $\\sim 10\\%$ of the orbit when the companion was at inferior conjunction. During the ingress and egress of the eclipse, delays in the pulse arrival times were observed which, along with the eclipse duration, suggested that the eclipses were due to excess ionized gas within the system (although the exact physical mechanism of the eclipses was and still is unclear eg. Thompson 1995). Optical studies showed large orbital variations in the light from the companion, suggesting the emission was primarily due to illumination by the pulsar (Fruchter et al. 1988b). H$\\alpha$ (Kulkarni et al. 1988) and X-ray (Stappers et al. 2003) imaging showed unambiguous evidence for a pulsar wind nebula around the pulsar, and observations of the non-thermal X-ray point source showed evidence for orbital modulation as well (eg. Huang and Becker 2007). Long term timing has shown that there are orbital period variations which are much larger in magnitude than would be expected from gravitational wave emission, with an orbital period derivative that can even change sign over time (Arzoumanian et al. 1994). These have  no easy interpretation, but mass loss from the system and an induced quadrapole moment of the companion could contribute to the measured changes in orbit (Applegate \\& Shaham 1994).  The inference drawn from these various observations is that PSR B1957+20 is ablating material off of its companion which shocks with the pulsar wind. Such systems had  been proposed as a stage in the formation of isolated millisecond pulsars by Ruderman et al. (1989). Since the pulsar seemed to be in the process of destroying its companion, it was dubbed the ``black widow\" pulsar. Early models of the shock emission expected from the original Black Widow pulsar were developed by Arons and Tavani (1993) and Raubenheimer et al (1995). These predicted electrons could be accelerated to $\\sim 3$ TeV and that orbital modulation could arise from obscuration by the companion of part of the shock, intrinsic beaming of the emission by the magnetic field in the shock, or by doppler boosting. The high energy luminosity could depend on the distance to the shock, the magnetic field of the pulsar, the ion fraction, the optical emission of the companion (in the case of inverse Compton emission), and the magnetization of the wind. However, the overall scaling should be primarily due to spin-down power of the pulsar $\\dot E$, the fraction of the wind involved in the shock $f$, and the distance to the system $d$.  \\begin{equation} L_S\\propto f \\dot E/d^2 \\end{equation}  No unambiguous detection of $\\gamma$-ray emission arising from the shock in PSR B1957+20 has yet been made (although there has been an intriguing suggestion of GeV emission in $Fermi$ data, see Wu et al. 2012). However, the ability of a pulsar wind to produce emission from radio frequencies up to TeV energies in an intrabinary shock has been demonstrated by observations of the PSR B1259$-$63 system (\\cite{jmm+99,aaa+11,hess13}) which consists of a young pulsar in an highly eccentric 3.4 year orbit around a main sequence Be star. Near periastron the radio pulses of PSR B1259$-$63 are obscured as it enters the region of the dense equatorial outflow of the Be star, and unpulsed radio, X-ray, GeV and TeV emission is observed. While the $\\dot E$ of PSR B1259$-$63 is somewhat higher and a larger fraction of the pulsar wind is shocked by the stellar wind than in PSR B1957+20, the total shock luminosity should not be much more than an order of magnitude greater. However, near periastron, PSR B1259$-$63 is a few hundred times brighter in X-rays than PSR B1957+20, despite being at similar distances.  The big difference between the Black Widow shock and those in PSR B1259-63 and the termination shocks of young pulsar wind nebulae is the scale of the systems. In terms of light cylinder radii $R_{lc}$, the companion of the Black Widow pulsar is on the order of $10^4 R_{lc}$ rather than the more typical $10^8 R_{lc}$ to the inner torii of young pulsar wind nebulae. This forces the shock to be in the region where, if the pulsar wind is magnetically dominated at the light cylinder, the magnetization fraction $\\sigma$ should still be relatively high. This may allow direct observational tests of proposed solutions to the ``sigma problem\" of pulsar wind shocks where the inferred magnetization at the termination shock from models of the X-ray emission is very low despite the expectation that at the light cylinder it should be very high (\\cite{a12}). Other potential advantages of studying the shocks in black widow type systems are that pulsar timing and optical observations allow us to determine the geometry of the system very well, and short orbits ($< 1$~d) allow for many repeated observations of the shock as viewed from different orientations.  For many years it seemed that eclipsing systems in the Galactic field might be quite rare. In the 20 years following the discovery of PSR B1957+20, only 2 other systems were discovered, both of which had spin-down fluxes $\\dot E/d^2$ much smaller than that of the original Black Widow (see Roberts 2011 and references therein). Observations of globular clusters, however, revealed that a significant fraction of their millisecond pulsars either eclipsed and/or had very low companion masses. Freire (2005) listed 18 such systems. 11 have very low mass companions ($M_c \\la 0.04 M_{\\odot}$)like PSR B1957+20, but 7 have masses more typical of helium white dwarfs. Optical observations of these latter systems tended to indicate the companions were, in fact, non-degenerate. Were these then systems where the pulsar was spun up by one companion, then exchanged the original companion for a non-degenerate one in the very crowded core of the globular cluster? If this were the case, then it would be natural that most such systems would be found in globular clusters. ",
        "conclusions": ""
    },
    "1402/1402.4737_arXiv.txt": {
        "abstract": "{ Chameleon, environmentally dependent dilaton, and symmetron gravity are three models of modified gravity in which the effects of the additional scalar degree of freedom are screened in dense environments. They have been extensively studied in laboratory, cosmological, and astrophysical contexts. In this paper, we  present a preliminary investigation into whether additional constraints can be provided by studying these scalar fields around black holes.  By looking at the properties of a static, spherically symmetric black hole,  we find that the presence of a non-uniform matter distribution induces a non-constant scalar profile in chameleon and dilaton, but not necessarily symmetron gravity.  An order of magnitude estimate shows that the effects of these profiles on in-falling test particles will be sub-leading compared to gravitational waves and hence observationally challenging to detect. } ",
        "introduction": "Since it was first published nearly a century ago, general relativity (GR) has earned its place as an incredibly successful and well verified theory of gravity (see e.g.\\ \\cite{Will2005}). There are however, both theoretical and observational reasons to consider alternatives to, and extensions of, GR. On the one hand, string theory provides a framework for describing quantum gravity, but suggests that we live in more than 4 dimensions. The consequences of these extra dimensions have not been definitively predicted, however, a generic feature of dimensional reduction is that extra fields appear in the low energy gravitational sector. On the other hand, observations of high redshift supernovae indicate that the universe is expanding at an accelerating rate~\\cite{Riess1998,Perlmutter1999}, and together with microwave background~\\cite{Hinshaw:2012aka,Ade:2013zuv} and large scale structure~\\cite{Abazajian:2008wr} measurements, suggest that around $70 \\%$ of the energy density of the universe comes in the form of a `dark energy' -- an energy momentum density which has a large negative pressure and is well modelled by a cosmological constant. Since the magnitude of the cosmological constant is extremely small by particle physics standards, explaining its stability under quantum corrections is a challenge. Finding either a natural explanation for its measured value or an alternative to the constant, is a major motivation for developing and studying modified theories of gravity (see \\cite{Clifton2011} for a review of various approaches). Additionally, studying modified gravity theories allows us to better explore where GR has been tested rigorously and to constrain the low energy properties of quantum gravity theories.  Scalar-tensor theories, which modify gravity by introducing new, non-minimally coupled scalar fields are an extensively studied alternative to GR. They are theoretically attractive because such light scalar fields are generically predicted in the low energy limit of string theory. For a scalar field to affect cosmological expansion, its mass must be of the order of the Hubble scale, $H_0\\sim 10^{-33}\\,\\text{eV}$.  If it interacts with matter however, the presence of a light scalar field will result in a long-range fifth force and would thus be subject to tight constraints from laboratory and solar system tests of gravity, \\cite{Will2005,Bertotti2003,Williams2004}. These constraints are weakened for theories in which the scalar field modifying gravity somehow decouples from matter, or is ``screened'', in dense environments. Models with this property are appealing because they can modify the behaviour of gravity on large scales while recovering the behaviour of GR in environments where local tests have been performed. This paper will focus on three particular scalar-tensor theories which have been found to exhibit screening, known as chameleon, \\cite{Khoury2004}, environmentally dependent dilaton, \\cite{Brax2010}, and symmetron, \\cite{Hinterbichler2011}, modified gravity.  These models have been studied and constrained using laboratory \\cite{Mota2008,Upadhye2012b,Upadhye2012,Upadhye2012a}, solar system~\\cite{Bertotti2003}, cosmological~\\cite{Jain2010}, and astrophysical~\\cite{Jain2012} tests. These investigations all have one thing in common: they probe gravity in a regime where gravitational fields and space-time curvatures are relatively weak.  In the near future, the direct detection of gravitational waves from compact binary systems will allow us to constrain the behaviour of gravity in the strong field, large curvature regime. Accordingly, attention has increasingly focussed on efforts to test gravity by studying the dynamics of compact objects such as neutron stars and black holes, \\cite{Yunes2013,Psaltis2008}. It is thus natural to ask whether observations of black holes might provide new constraints on screened modified gravity.  Studying black holes in the context of a modified gravity theories inevitably leads to the re-examination of the uniqueness of exact solutions, i.e. ,whether the extra physical fields add extra degrees of freedom to black hole solutions, usually referred to as ``black hole hair'', \\cite{Ruffini}. A slightly looser definition of ``hair'', in terms of non-trivial scalar profiles rather than extra measurable charges at infinity known as {\\it dressing} \\cite{Achucarro:1995nu}, has been known to be possible for many years, starting with the (unstable) coloured black holes \\cite{Bartnik:1988am}, or the (stable) black monopoles, \\cite{Lee:1991vy}, quantum hair, \\cite{Coleman:1991jf,Dowker:1991qe}, as well as the explicit vortex hair \\cite{Achucarro:1995nu}.  In scalar tensor gravity, a number of  no-hair theorems (i.e., demonstrating that the scalar fields takes a constant value around an isolated black hole) have been proven, \\cite{Bekenstein1996,Sotiriou2012,Faraoni2013} (although note the assumptions behind these theorems may not always correspond to desired physical situations, \\cite{Sotiriou:2013qea}).  The results of these theorems have been extended to binary black hole systems using perturbative, \\cite{Yunes2012,Mirshekari2013}, and numerical, \\cite{Healy2011}, calculations, demonstrating that dynamical spacetimes with two interacting black holes in scalar-tensor theory will be indistinguishable from general relativity. This might seem to imply that black hole systems will not be useful for constraining screened modified gravity, however, these theorems do not take into account cosmological backgrounds, in which the scalar field will typically be dynamical. A study of cosmologically evolving black holes with a canonical scalar field, \\cite{Chadburn:2013mta}, shows that the scalar field does evolve on the event horizon of the black hole, though there is no evidence for an additional scalar charge. Indeed, there are several solutions in the literature which allow for nontrivial scalar fields around black holes, \\cite{Fonarev:1994xq,Clifton:2004st,Clifton:2006ug,Guariento:2012ri}, although to be fair, many are singular, or ``engineered''.  This paper focuses on removing the requirement that the black hole exists in a vacuum.  This consideration is clearly relevant for astrophysical black holes that are typically observed in galaxies or galactic centres, and can have energetic accretion discs. In order to explore this question, we consider spherically symmetric distributions of matter around a black hole in screened modified gravity. We perform a `probe' calculation, in that we explore the scalar profile around the black hole without considering the corresponding modification of the black hole geometry. We are therefore not looking at issues of time dependent black holes, such as discussed in \\cite{McVittie:1933zz,Sultana:2005tp,Faraoni:2007es,Carrera:2008pi,Abdalla:2013ara}, as we will see nontrivial profiles even in the static case.  Strictly speaking, a non-uniform matter distribution will result in a non-uniform scalar field profile due to the non-minimal coupling of the scalar to gravity. We expect that the complete picture will be a superposition of multipoles, with the dominant features being encapsulated by the monopole, or spherically symmetric, behaviour. Indeed, if we were to discover that the spherically symmetric case precluded the possibility of black hole hair, then it would be indicative that black holes would not carry hair. Conversely, the discovery of non-trivial spherically symmetric scalar profiles would demonstrate that black holes can indeed carry screened scalar hair, although the full solution would be more complicated and involved than the simple picture presented here.  The layout of the paper is as follows: we introduce screened modified gravity in \\textsection\\ref{screenedMG}. We then study the effect of matter on the scalar field profile in \\textsection\\ref{profile} under the assumption of spherical symmetry and discuss observational implications for black hole properties in \\textsection\\ref{implications}. ",
        "conclusions": "The strong field, large curvature limit of modified gravity is a largely unexplored area for constraining various model dependent parameters and differentiating between the numerous models that exist in literature. We have presented an exploratory calculation for theories of modified gravity with screening mechanisms by studying a static, spherically symmetric black hole with an $r$-dependent matter distribution around it. In this construction, we found that chameleon and dilaton fields develop a non-trivial scalar profile, while the symmetron assumes a constant value. An order of magnitude estimate showed that the resulting scalar gradients affect in falling test particles in a way that is sub-leading compared to the quadruple radiation in GR and thus would be observationally challenging to detect.  Note that our findings for symmetrons could be qualitatively altered if we used a different matter distribution in our setup. As we discussed in \\textsection\\ref{sym_profile}, if the matter density far from the black hole is small enough to make the symmetron approach $\\phistar=\\mu/\\sqrt{\\lambda}$ as its boundary condition, its field equation will no longer have a constant solution. This could occur, for example, if we used a model of an accretion disk which had finite spatial extent.  We would, however, expect any symmetron gradients to be comparable to the ones found above for long range chameleons and dilatons.  Similar reasoning can be applied when we consider solutions to the scalar field equation with angular dependence, or solutions on a Kerr background.   No hair theorems  require that in a stationary system involving matter around a black hole, any gradient in the scalar field will be sourced entirely by non-uniformities in matter density.  If we assume that the density contrasts examined above are typical of astrophysical matter distributions near black holes, the magnitude of the scalar gradients produced (and their consequent fifth forces) should likewise be representative.  Thus generically we would expect theories of screened modified gravity to produce solutions distinct from GR around stationary black hole systems with sub leading corrections to various observational effects.  It would be interesting to explore this strong field limit further. One possible physical effect to study within our setup would be the phase of gravitational waves. Our crude estimates suggest that the rate of GR quadruple radiation will dominate over any radiation in the scalar sector, however,  in~\\cite{Berti2012}, the authors showed that in Brans-Dicke theory the presence of a massive scalar field affects the phasing of gravitational radiation from binary systems significantly and they found that for observations of  intermediate mass ratio inspirals, this effect could be used to place constraints on model parameters (a lower bound on $\\omega_{BD}$ and an upper bound on the mass of the scalar) which are competitive with Cassini and LLR measurements.  Additionally, the requirement that our system and solutions be stationary causes us to disregard the possibility of transient effects associated with superradiance.  Superradiance is a property of rotating black holes through which incident waves with at resonant frequency can become amplified by extracting some of the black hole's rotational energy. A number of studies have shown that  superradiant effects can give rise to long-lived unstable modes in the presence of a massive scalar \\cite{Cardoso2005,Dolan2007,Dolan2012,Witek2013,Cardoso2013a}. This occurs for isolated black holes, but it also has been shown that when a black hole is surrounded by matter, superradiant modes can be amplified by factors as large as $10^5$~\\cite{Cardoso2013}. If the scalar's mass is very light, the instability timescale for these modes becomes short and would result in gaps in the mass-spin phase space of observed black holes. Because of this, measurements of a black hole's mass and spin can be used to constrain allowed masses for scalar fields.  This technique has already been used to place the most stringent upper bound on the mass of the photon~\\cite{Cardoso2013a,Berti2013a}.  Yet another potentially observable effect of superradiance occurs in EMRI systems, in what is known as floating orbits \\cite{Yunes2013,Yunes2012,Cardoso2011}.  Floating orbits occur when the orbital frequency of the small compact object can excite superradiant modes.  This causes the scalar field to transfer rotational energy from the central black hole into the orbit of the small compact object, counteracting the energy lost through quadrupole radiation and slowing the inspiral. Such floating orbits would affect the EMRI's gravitational waveforms significantly, so the detection of an EMRI signal consistent with GR would allow us to place strong constraints on the mass of a light scalar.  We note with interest that these effects are typically studied in the context of scalar fields with a mass in the range $10^{-33}\\,\\text{eV} - 10^{-10}\\,\\text{eV}$~\\cite{Berti2013a}, which is relevant for dilatons and symmetrons, as well as chameleons at galactic or cosmological densities.  Consideration of gravitational wave phasing, superradiant instabilities, and floating orbits will thus be important if we want to understand how these theories can be constrained by observations of black hole systems. We hope to further explore some of these effects in future works."
    },
    "1402/1402.6442_arXiv.txt": {
        "abstract": "We present the association rates between solar energetic particles (SEPs) and the radio emission signatures in the corona and IP space during the entire solar cycle 23. We selected SEPs associated with X and M-class flares from the visible solar hemisphere. All SEP events are also accompanied by coronal mass ejections. Here, we focus on the correlation between the SEP events and the appearance of radio type II, III and IV bursts on dynamic spectra. For this we used the available radio data from ground-based stations and the Wind/WAVES spacecraft. The associations are presented separately for SEP events accompanying activity in the eastern and western solar hemisphere. We find the highest association rate of SEP events to be with type III bursts, followed by types II and IV. Whereas for types III and IV no longitudinal dependence is noticed, these is a tendency for a higher SEP-association rate with type II bursts in the eastern hemisphere. A comparison with reports from previous studies is briefly discussed. ",
        "introduction": "Solar energetic particle (SEP) events are transient flux enhancements of electrons, protons and ions due to acceleration processes in the solar corona and interplanetary (IP) space. The high energy particles can pose a serious risk for the near Earth and ground-based technological devices, may disturb communications and be a health hazard \\citep{2006SSRv..124..303M}. This is why energetic particles are of major space weather interest. The accelerator of these particles, however, is still a subject of debate. Presently, the reconnection processes during flares \\citep{2002JGRA..107.1315C} and the shock-acceleration at coronal mass ejections (CMEs), \\cite{1999SSRv...90..413R}, are recognized as efficient particle accelerators. An important issue is whether both acceleration processes act simultaneously with comparable efficiency or one of them dominates the particle energization process. In the current two-class SEP picture, the flares are thought to dominate the {\\it impulsive} events and the shock waves dominate the large {\\it gradual} events. Observations, however, are not able to relate unambiguously the SEP parameters measured in situ to the parent solar activity (e.g., flares vs. shocks). The large uncertainties usually present when doing different correlation analyses may arise due to the poorly known particle transport in the turbulent IP magnetic field, the magnetic connection between the acceleration site and the Earth, and that SEPs are often just measured at a single point in space.  Radio observations can provide additional information and constraints for identifying the particle accelerator in the corona and IP space. When nonthermal electrons propagate through the solar corona they may emit radio waves, if certain conditions apply \\citep{2008SoPh..253....3N}. The different types of radio bursts in the dynamic radio spectrum (see e.g., a review by \\citet{2008A&ARv..16....1P}), are usually interpreted as signatures from electrons accelerated in the vicinity of a shock wave (type IIs), as electron beams (type IIIs) or as electrons confined in closed loop structure during the early evolution of a CME (type IVs). Type III bursts usually extend to low frequencies (reaching as far as the Earth orbit, $\\sim$20 kHz) which is the main indicator that accelerated electrons (and by inference also protons) can efficiently escape from the solar corona \\citep{2002JGRA..107.1315C}. Interplanetary counterparts of type II are also observed, but to lesser extent, whereas type IVs can be seen only to few MHz, but not to lower frequencies. Here we focus on the comparison of the particle (mainly proton) intensities measured in situ with the electromagnetic emission of electrons in the corona and IP space over the entire solar cycle 23.  Numerous previous studies reported different association rates of SEPs and emission of radio bursts of type II \\citep{2004ApJ...605..902C,2008AnGeo..26.3033G}, III \\citep{2002JGRA..107.1315C} and IV \\citep{1982ApJ...261..710K}. The aim here is to identify the different types of radio emission from the dynamic spectra, to complement our results with reports from the different radio observatories and to compare them with previous works on the topic. ",
        "conclusions": "We present the association rates of the SEP events (protons) and their accompanying radio emission (from electrons) in the corona (dm and m wavelength) and IP space (DH-range). Since there are no signatures of protons interacting with the solar atmosphere (with the exception of gamma-ray emission), we use electron signatures as a diagnostic for particle acceleration from the corona up to 1~AU.  \\cite{2002JGRA..107.1315C} were the first to identify long-lasting groups of DH type III bursts as a typical radio counterpart of large SEP events. They found the groups to be of significantly longer duration than type~III bursts associated with impulsive flares, see also \\cite{1987SoPh..111..397M}. In the analysis performed here, we did not take into account the burst duration of DH-type III, nor explicitly separate the SEP events into gradual or impulsive. Still we find that SEP events have the highest association rate with type III radio bursts. This implies that the electrons accelerated in the corona (within one solar radius) have a ready access to the IP space, irrespective of whether the SEP event is impulsive or gradual.  Shock signatures (type II bursts) were also considered in correlation studies with SEP events. We found a lower association rate (up to 75\\%) of the DH-type II bursts with SEP events, compared to the m- and DH-type III association rates. The number increases when only SEP events with strong intensities are considered, in agreement with \\cite{2002ApJ...572L.103G} and \\cite{2004ApJ...605..902C}.  The association of SEP events with types III and IV is comparable in the eastern and western groups (see Figure~\\ref{F-histo}). But for the type II bursts there is a slight trend for a higher association rate in the eastern hemisphere. This result could be understood it terms of different sources contributing to SEP events. Shocks could be the dominant accelerator when the parent activity is poorly connected to Earth (as in the eastern solar hemisphere). In the western hemisphere (where more than twice as many events were detected) both flare and shock acceleration could contribute. When no shock signatures accompany the eastern SEP events, their propagation through the IP space and detection at Earth could be facilitated by a large-scale magnetic structure, e.g., interplanetary coronal mass ejections (ICMEs). \\cite{1991JGR....96.7853R} estimated that for about 15\\% of the eastern events this is likely the case. Recently \\cite{2013SoPh..282..579M} showed similar occurrence rate of ICMEs for the western SEP events (20\\%).  The association with the type IIIs seems to be increasing with the decreasing of the radio frequency (from dm to DH-range). At first, the type III identification in the dm-range might be masked due to overlying decimeter continuum often present in the dynamic radio spectra (observational bias). Hence, the obtained association rates of dm-IIIs are to be considered as lower limits only. In addition, high frequencies and dense plasma impede the production and growth of Langmuir waves. Also, the electron beam must travel an `instability distance' from the acceleration site before Langmuir waves are generated \\citep{2011A&A...529A..66R}. This can cause a lack of high frequency type III emission in even the most energetic of electron beams. Moreover, collisions of both waves and electrons can suppress Langmuir wave growth at high frequencies, e.g., \\citet{1982ApJ...263..423K,2012SoPh..tmp..109R}.  Without the intention to give a complete overview on the results from previous work, we selected few statistical studies that are (partially) co\\-vering solar cycle 23. The different association rates are given in Table~\\ref{T-Rates_Sum} and are mostly consistent, although in many earlier studies only large SEP events were considered."
    },
    "1402/1402.3267_arXiv.txt": {
        "abstract": "{ Scaling properties of galaxy cluster observables with cluster mass provide central insights into the processes shaping clusters.  Calibrating proxies for cluster mass that are relatively cheap to observe  will moreover be crucial to harvest the cosmological information available from the number and growth of clusters  with upcoming surveys like \\textit{eROSITA} and \\textit{Euclid}. The recent \\textit{Planck} results led to suggestions that X-ray masses might be biased low by $\\sim\\!40$~\\%, more than previously considered. }{ We aim to extend  knowledge of the weak lensing -- X-ray mass scaling towards lower masses (as low as $1\\!\\times\\!10^{14}\\,\\mathrm{M}_{\\odot}$) in a sample representative of the $z\\!\\sim\\!0.4$--$0.5$ population. Thus, we extend the direct calibration of cluster mass estimates to higher redshifts. }{ We investigate the scaling behaviour of MMT/Megacam weak lensing (WL) masses for $8$ clusters at $0.39\\!\\leq\\!z\\!\\leq\\!0.80$ as part of the \\emph{400d} WL programme with hydrostatic \\textit{Chandra} X-ray masses as well as those based on the proxies, e.g. $Y_{\\mathrm{X}}\\!=\\!T_{\\mathrm{X}}M_{\\mathrm{gas}}$. }{ Overall, we find good agreement between WL and X-ray masses, with different mass bias estimators all consistent with zero. When subdividing the sample into a low-mass and a high-mass subsample, we find the high-mass subsample to show no significant mass bias while for the low-mass subsample, there is a bias towards overestimated X-ray masses at the $\\sim\\!2\\sigma$ level for some mass proxies. The overall scatter in the mass-mass scaling relations is surprisingly low. Investigating possible causes, we find that neither the greater range in WL than in X-ray masses nor the small scatter can be traced back to the parameter settings in the WL analysis. }{ We do not find evidence for a strong ($\\sim\\!40$~\\%) underestimate in the X-ray masses, as suggested to reconcile recent \\textit{Planck} cluster counts and cosmological constraints. For high-mass clusters, our measurements are consistent with other studies in the literature. The mass dependent bias, significant at $\\sim\\!2\\sigma$, may hint at a physically different cluster population (less relaxed clusters with more substructure and mergers); or it may be due to small number statistics.  Further studies of low-mass high-$z$ lensing clusters will elucidate their mass scaling behaviour. } ",
        "introduction": "Galaxy cluster masses hold a crucial role in cosmology. In the paradigm of hierarchical structure formation from tiny fluctuations in the highly homogeneous early cosmos after inflation, clusters emerge via the continuous matter accretion onto local minima of the gravitational potential. Depending sensitively on cosmological parameters, the cluster mass function, i.e.\\ their abundance as function of mass and redshift $z$, provides observational constraints to cosmology \\citep[e.g.,][]{2009ApJ...692.1060V,2011ARA&A..49..409A,2013arXiv1303.5080P}.  Observers use  several avenues to determine cluster masses: properties of the X-ray--emitting intracluster medium (ICM), its imprint on the cosmic microwave background via the Sunyaev-Zel'dovich (SZ) effect in the sub-mm regime, galaxy richness estimates and dynamical masses via optical imaging and spectroscopy, and gravitational lensing. Across all wavelengths, cluster cosmology surveys are under preparation, aiming at a complete cluster census  out to ever higher redshifts, e.g. \\textit{eROSITA} \\citep{2010SPIE.7732E..23P,2012arXiv1209.3114M,2012MNRAS.422...44P} and \\textit{Athena} \\citep{2013arXiv1306.2307N,2013arXiv1306.2319P} in X-rays, \\textit{Euclid} \\citep{2011arXiv1110.3193L,2012arXiv1206.1225A}, DES and LSST in the optical/near-infrared, CCAT \\citep{2012SPIE.8444E..2MW} and SKA at sub-mm and radio frequencies.  Careful X-ray studies of  clusters at  low and intermediate redshifts  yield highly precise cluster masses, but assume hydrostatic equilibrium, and in most cases spherical symmetry \\citep[e.g.][]{2008A&A...487..431C,2013SSRv..177..119E}. Observational evidence and numerical modelling challenge these assumptions for all but the most relaxed systems \\citep[e.g.][]{2008MNRAS.384.1567M,2012NJPh...14e5018R,2013SSRv..177..155L,2013ApJ...765...24N}. While  simulations find X-ray masses to only slightly underestimate the true mass of clusters that exhibit no indications of recent mergers and can be considered virialised, non-thermal pressure support can lead to a $>\\!20$~\\% bias in unrelaxed clusters \\citep{2010A&A...510A..76L,2012NJPh...14e5018R}. \\citet{2014ArXiv1401.7657S} modelled the pressure due to ICM turbulence analytically and found a $\\sim\\!10$~\\% underestimate of cluster masses compared to the hydrostatic case.  Weak lensing (WL), in contrast, is subject to larger stochastic uncertainties, but can in principle yield unbiased masses, as no equilibrium assumptions are required. Details of the mass modelling however can introduce biases, in particular concerning projection effects, the source redshift distribution and the departures from an axisymmetric mass profile \\citep{2009MNRAS.396..315C,2011ApJ...740...25B,2012MNRAS.421.1073B,2013SSRv..177...75H}. For individual clusters, stochastic uncertainties dominate the budget; however, larger cluster samples benefit from improved corrections for  lensing systematics, driven by cosmic shear projects \\citep[e.g.][]{2013MNRAS.429..661M}.  Most of the leverage on cosmology and structure formation from future cluster surveys will be due to clusters at higher $z$ than have been previously investigated. Hence, the average cluster masses and signal/noise ratios for all observables are going to be smaller. Even and especially for the deepest surveys, most objects will lie close to the detection limit. Thus the scaling of inexpensive proxies (e.g. X-ray luminosity $L_{\\mathrm{X}}$) with total mass needs to be calibrated against representative cluster samples at low and high $z$. Weak lensing and SZ mass estimates are both good candidates as they exhibit independent systematics from X-rays and a weaker $z$-dependence in their signal/noise ratios.  Theoretically, cluster scaling relations arise from their description as self-similar objects forming through gravitational collapse \\citep{1986MNRAS.222..323K}, and deviations from simple scaling laws provide crucial insights into cluster physics. For the current state of scaling relation science, we point to the recent review by \\citet{2013SSRv..177..247G}. As we are interested in the cluster population to be seen by upcoming surveys, we focus here on results obtained at high redshifts.  Self-similar modelling includes evolution of the scaling relation \\emph{normalisations} with the Hubble expansion, which is routinely measured \\citep[e.g.][]{2011A&A...535A...4R, 2013AN....334..354E}. Evolution effects beyond self-similarity, e.g.\\ due to declining AGN feedback at low $z$, have been claimed and discussed \\citep[e.g.][]{2007MNRAS.382.1289P,2010MNRAS.408.2213S,2010ApJ...715.1508S,2012MNRAS.421.1583M}, but current observations are insufficient to constrain possible evolution in slopes \\citep{2013SSRv..177..247G}. Evidence for different scaling behaviour in groups and low-mass clusters was found by, e.g.,\\ \\citet{2011A&A...535A.105E,2012MNRAS.422.2213S,2014arXiv1402.0868B}.  \\citet{2011A&A...535A...4R} and \\citet{2012MNRAS.421.1583M} investigated X-ray scaling relationships including clusters at $z\\!>\\!1$, and both stressed the increasing influence of selection effects at higher $z$. Larger weak lensing samples of distant clusters are just in the process of being compiled \\citep{2011ApJ...737...59J,2012A&A...546A.106F,2011ApJ...726...48H,2012MNRAS.427.1298H,2012A&A...546A..79I,2012arXiv1208.0597V,2012ApJS..199...25P}. Thus most WL scaling studies are currently limited to $z\\!\\!\\lesssim\\!\\!0.6$, and also include nearby clusters \\citep[e.g.][M13]{2012MNRAS.427.1298H,2013ApJ...767..116M}. The latter authors find projected WL masses follow the expected correlation with the  SZ signal $Y_{\\mathrm{SZ}}$, corroborating similar results for more local clusters by \\citet{2009ApJ...701L.114M,2012ApJ...754..119M}. \\citet{2013MNRAS.429.3627M} performed a detailed WL analysis of a $z\\!=\\!0.81$ cluster discovered in the SZ using the Atacama Cosmology Telescope, and compared the resulting lensing mass against the \\citet{2012ApJ...751...12R} $Y_{\\mathrm{SZ}}$--$M$ scaling relation, in what they describe as a first step towards a high-$z$ SZ-WL scaling study.  By compiling \\textit{Hubble Space Telescope} data for $27$ massive clusters at $0.83\\!<\\!z\\!<\\!1.46$, \\citet{2011ApJ...737...59J} not only derive the relation between WL masses $M^{\\mathrm{wl}}$ and ICM temperature $T_{\\mathrm{X}}$, but also notice a good correspondence between WL and hydrostatic X-ray masses $M^{\\mathrm{hyd}}$. As they focus on directly testing cosmology with the most massive clusters , these authors however stop short of deriving the WL--X-ray scaling. Also using \\textit{HST} observations, \\citet{2011ApJ...726...48H} investigated the WL mass scaling of the optical cluster richness (i.e.\\ galaxy counts) and $L_{\\mathrm{X}}$ of $25$ moderate-$L_{\\mathrm{X}}$ clusters at $0.3\\!<\\!z\\!<\\!0.6$, thus initiating the study of WL scaling relations off the top of the mass function.  Comparisons between weak lensing and X-ray masses for larger cluster samples were pioneered by \\citet{2008MNRAS.384.1567M} and \\citet{2008A&A...482..451Z}, collecting evidence for the ratio  of weak lensing  to X-ray masses $M^{\\mathrm{wl}}/M^{\\mathrm{hyd}}\\!>\\!1$, indicating non-thermal pressure. \\citet{2010ApJ...711.1033Z}, analysing $12$ clusters from the \\textit{Local Cluster Substructure Survey} (LoCuSS), find this ratio to depend on the radius . Likewise, a difference between relaxed and unrelaxed clusters is found \\citep{2010ApJ...711.1033Z,2013ApJ...767..116M}. \\citet{2012NJPh...14e5018R} show that the gap between X-ray and lensing masses is more pronounced in simulations than in observations, pointing to either an underestimate of the true mass also by WL masses \\citep[cf.][]{2012MNRAS.421.1073B} or to simulations overestimating the X-ray mass bias.  The current disagreement between the cosmological constraints derived from \\textit{Planck} primary cosmic microwave background (CMB) data with \\textit{Wilkinson Microwave Anisotropy Probe} data, supernova data, and cluster data \\citep{2013arXiv1303.5080P} may well be alleviated by, e.g.\\ sliding up a bit along the \\textit{Planck} degeneracy curve between the Hubble factor $H_{0}$ and the matter density parameter $\\Omega_{\\mathrm{m}}$. Nevertheless, as stronger cluster mass biases than currently favoured $(\\sim\\!40$~\\%) have also been invoked as a possible explanation, it is very important to test the cluster mass calibration with independent methods out to high $z$, as we do in this work.  This article aims to test the agreement of the weak lensing and X-ray masses measured by \\citet{2012A&A...546A..79I} for $8$ relatively low-mass clusters at $z\\!\\gtrsim\\!0.4$ with scaling relations from the recent literature. The \\emph{400d} X-ray sample from which our clusters are drawn has been constructed to contain typical objects at intermediate redshifts, similar in mass and redshift to upcoming surveys. Hence, it does not include extremely massive low-$z$ clusters. We describe the observations and WL and X-ray mass measurements for the $8$ clusters in Sect.~\\ref{sec:obsdat}, before presenting the central scaling relations in Sect.~\\ref{sec:res}. Possible explanations for the steep slopes our scaling relations exhibit are discussed in Sect.~\\ref{sec:puzzle}, and we compare to literature results in Sect.~\\ref{sec:m13disc}, leading to the conclusions in Sect.~\\ref{sec:conclu}. Throughout this article, $E(z)\\!=\\!H(z)/H_{0}\\!=\\!\\sqrt{\\Omega_{\\mathrm{m}}(z+1)^{3}+\\Omega_{\\Lambda}}$ denotes the self-similar evolution factor (Hubble factor $H(z)$ normalised to its present-day value of $H_{0}\\!=\\!72\\,\\mbox{km}\\,\\mbox{s}^{-1}\\,\\mbox{Mpc}^{-1}$), computed for a flat universe with matter and Dark Energy densities of $\\Omega_{\\mathrm{m}}\\!=\\!0.3$ and $\\Omega_{\\Lambda}\\!=\\!0.7$ in units of the critical density. ",
        "conclusions": "\\label{sec:conclu}  In this article, we analysed the scaling  relation between WL and X-ray masses for $8$ galaxy clusters drawn from the \\emph{400d} sample of X-ray--luminous $0.35\\!\\leq\\!z\\!\\leq\\!0.89$ clusters. WL masses were measured from the \\citet{2012A&A...546A..79I} MMT/Megacam data, and X-ray masses were  based on the V09a \\textit{Chandra} analysis. We summarise our main results as follows:  \\begin{description} \\item[1.] We probe the WL--X-ray mass scaling relation, in an unexplored region of the parameter space for the first time: the $z\\!\\sim\\!0.4$--$0.5$ redshift range, down to $1\\!\\times\\!10^{14}\\,\\mathrm{M}_{\\odot}$. \\item[2.] Using several X-ray mass estimates, we find the WL and X-ray masses to be consistent with each other. Most of our clusters are compatible with the $M^{\\mathrm{X}}\\!=\\!M^{\\mathrm{wl}}$ line. \\item[3.] Assuming the $M^{\\mathrm{wl}}$ not to be significantly biased, we do not find evidence for a systematic underestimation of the X-ray masses by $\\sim\\!40$~\\%, as suggested as a possible solution to the discrepancy between the \\textit{Planck} CMB constraints on $\\Omega_{\\mathrm{m}}$ and $\\sigma_{8}$ (the normalisation of the matter power spectrum) and the \\textit{Planck} SZ cluster counts \\citep{2013arXiv1303.5080P}. While our results favour a small WL--X-ray mass bias, they are consistent with both vanishing bias and the $\\sim\\!20$~\\% favoured by studies of non-thermal pressure support. \\item[4.] For the mass-mass scaling relations involving  $M^{\\mathrm{wl}}$, we observe a surprisingly low scatter $0.5\\!\\!<\\!\\!\\chi^{2}_{\\mathrm{red}}\\!\\!<\\!\\!0.6$, although we use only stochastic uncertainties and allow for correlated errors via a Monte Carlo method. Because the errors in $M^{\\mathrm{wl}}$ are largely determined by the intrinsic WL shape noise $\\sigma_{\\varepsilon}$, we however deem a drastic overestimation unlikely (Sect.~\\ref{sec:chisqdisc}). For the scaling relations involving $M_{\\mathrm{G}}$, however, we observe a large scatter, contrary to \\citet{2010ApJ...721..875O} and M13. \\item[5.] Looking in detail, there are intriguing indications for a mass-dependence of the WL--X-ray mass ratios of our relatively low-mass $z\\!\\sim\\!0.4$--$0.5$ clusters. We observe a mass bias in the low--$M^{\\mathrm{wl}}$ mass bin at the $\\sim\\!2\\sigma$ level when splitting the sample at $\\log{(M_{\\mathrm{piv}}/\\mathrm{M}_{\\odot})}\\!=\\!14.5$ This holds for the masses V09a report based on the $Y_{\\mathrm{X}}$, $T_{\\mathrm{X}}$, and $M_{\\mathrm{G}}$ proxies. \\end{description}  We thoroughly investigate possible causes for the mass-dependent bias and tight scaling relations. First (Sect.~\\ref{sec:c200}), we confirm that by using a mass-concentration relation instead of directly fitting $c_{200}$ from WL, we already significantly reduced the bias due to conversion from $r_{200}$ to $r_{500}$. We emphasise that, on average, the NFW shear profile represents a suitable fit for the cluster population \\citep[cf.][]{2013ApJ...769L..35O}. Measuring $M^{\\mathrm{hyd}}$ within $r_{500}^{\\mathrm{wl}}$ induces correlation between the data points in Fig.~\\ref{fig:mxml}. Removing this correlation by plotting both masses within a fixed physical radius, we still find small scatter (Sect.~\\ref{sec:chisqdisc}).  We notice that the mass range occupied by the $M^{\\mathrm{wl}}$ exceeds the X-ray mass ranges, Partially, this higher WL mass range can be explained by  the correction for dilution by member galaxies, which could be applied only where colour information was available (Paper~II). Coincidentally, this is the case for the more massive half of the MMT sample in terms of $M^{\\mathrm{wl}}$, thus boosting the range of measured WL masses (Sect.~\\ref{sec:ndc}). This result underscores the importance of correcting for the unavoidable inhomogeneities in WL data due to the demanding nature of WL observations \\citep[cf.][]{2012arXiv1208.0605A}. We find no further indications for biases via the WL analysis. Furthermore the tight scaling precludes strong redshift effects, and we find that our small MMT subsample is largely representative of the complete sample of $36$ clusters, judging from the $M^{\\mathrm{Y}}$--$M^{\\mathrm{T}}$ relation (Sect.~\\ref{sec:fluke}). For the $M^{\\mathrm{Y}}$--$M^{\\mathrm{G}}$ and $M^{\\mathrm{T}}$--$M^{\\mathrm{G}}$ relations, significant scatter ($\\chi^{2}_{\\mathrm{red}}\\!>\\!2$) is present in the larger sample. The former relation also shows indications for a significant bias of $M^{\\mathrm{Y}}\\!\\approx\\!1.15 M^{\\mathrm{G}}$.  Weak lensing and hydrostatic masses for the \\emph{400d} MMT clusters are in good agreement with the $z\\!>\\!0.35$ part of the \\citet{2013ApJ...767..116M} sample and the $M^{\\mathrm{wl}}_{500}$--$M^{\\mathrm{hyd}}_{500}$ relation derived from it (Sect.~\\ref{sec:hydlit}). The M13 and \\citet{2012A&A...546A.106F} samples include three \\emph{400d} clusters with CFHT WL masses. These clusters neither point to significantly higher scatter nor to a less mass-dependent bias (Fig.~\\ref{fig:lit}). We are planning a re-analysis of the CFHT data, having demonstrated in Paper~II that lensing catalogues from MMT and CFHT are nicely compatible. Such reanalysis is going to be helpful to identify more subtle WL analysis effects potentially responsible for the steep slopes and tight correlation of WL and X-ray masses.  An alternative explanation are intrinsic differences in the low-mass cluster population. That the \\emph{400d} MMT sample probes to slightly lower masses ($1\\!\\times\\!10^{14}\\,\\mathrm{M}_{\\odot}$) than M13 or F12 becomes especially obvious from the $M^{\\mathrm{wl}}$--$Y^{\\mathrm{X}}$ relation (Fig.~\\ref{fig:lit}, Sect.~\\ref{sec:yxlit}). Because the \\emph{400d} sample is  more representative of the $z\\!\\sim\\!0.4$--$0.5$ cluster population, it is likely to contain more significant mergers relative to the cluster mass, skewing mass estimates (Sect.~\\ref{sec:physics}). Hence, the \\emph{400d} survey might be the first to see the onset of a mass regime in which cluster physics and substructure lead the WL--X-ray scaling to deviate from what is known at higher masses. Remarkably, Giles et al.\\ (in prep.) are finding a different steep slopes in their low-mass WL--X-ray scaling analysis. Detailed investigations of how their environment shapes clusters like CL\\,1416+4446 might be necessary to improve our understanding of the cluster population to be seen by future cosmology surveys. Analysis systematics might also behave differently at lower masses. A turn for WL cluster science towards lower mass objects, e.g.\\ through the completion of the \\emph{400d} WL sample, will help addressing the question of evolution in lensing mass scaling relations."
    },
    "1402/1402.3857.txt": {
        "abstract": "Two formation scenarios have been proposed to explain the tight orbits of hot Jupiters. They could be formed in orbits with a small inclination (with respect to the stellar spin) via disk migration, or in more highly inclined orbits via high-eccentricity migration, where gravitational interactions with a companion and tidal dissipation are at play. Here we target hot Jupiter systems where the misalignment $\\lambda$ has been inferred observationally and we investigate whether their properties are consistent with high-eccentricity migration. Specifically, we study whether stellar tides can be responsible for the observed distribution of $\\lambda$ and orbital separations. Improving on previous studies, we use detailed models for each star, thus accounting for how convection (and tidal dissipation) depends on stellar properties. In line with observations suggesting that hotter stars have higher $\\lambda$, we find that $\\lambda$ increases as the amount of stellar surface convection decreases. This trend supports the hypothesis that tides are the mechanism shaping the observed distribution of $\\lambda$. Furthermore, we study the past orbital evolution of five representative systems, chosen to cover a variety of temperatures and misalignments. We consider various initial orbital configurations and integrate the equations describing the coupled evolution of the orbital separation, stellar spin, and misalignment.  We account for stellar tides and wind mass loss, stellar evolution, and magnetic braking. We show that the current properties of these five representative systems can be explained naturally, given our current understanding of tidal dissipation and with physically motivated assumptions for the effects driving the orbital evolution. ",
        "introduction": "\\label{Intro} The plethora of exoplanets discovered in recent years has revealed that planetary systems exist in a much greater variety than we had ever imagined. To date, more than 1000 exoplanets have been confirmed using different observational techniques. Almost 200 of these planets are similar in mass to Jupiter, but revolve around their parent star every 10\\,days or less (NASA Exoplanet Archive), thus challenging our understanding of planet formation and evolution. Different scenarios have been proposed to explain how these so-called {\\it hot Jupiters} formed in such tight orbits. % and an additional 3,600 candidates have been provided by NASA's {\\it Kepler}'s satellite. These discoveries have shown that planetary systems exist in a much greater variety than we had ever imagined. In this paper we consider hot Jupiters, giant planets which challenge our understanding of planet formation and evolution by revolving around their parent stars in very tight orbits. Different formation scenarios have been proposed to explain these planets. One way to distinguish between these models is to investigate the current properties of the many discovered systems. Our focus here lies on systems where the obliquity $\\lambda$ (the sky-projected angle between the stellar spin and orbital angular momentum vectors) has been constrained observationally.  Two migration models have been invoked to bring gas giants from their birthplace at several AU into the tight orbits we observe today: {\\it disk migration} and {\\it high-eccentricity migration} (however, see also \\citealt{TutukovFedorova2012, Thies+11}). These models predict different orientations of the planet's orbit at present. In the disk migration scenario, planets could migrate inwards through their interactions with the protoplanetary gas disk \\citep{GoldreichTremaine80,Ward97, MurrayHHT98, Lin+96,Guillochon+11}. As the disks tend to damp the orbital inclination \\citep{Cresswell+07,XiangPapaloizou14}, this model would naturally lead to nearly circular orbits and small obliquities (e.g., \\citealt{GoldreichTremaine80,PapaloizouLarwood00}). In the high-eccentricity migration scenario, gravitational interactions either between several planets or with companion stars could lead to highly eccentric orbits and high obliquities (\\citealt{Kozai62, Lidov62, WuLithwick11,Naoz+11,Nagasawa08,FabryckyTremaine07,WuMurray03,RasioFord96, ChatterjeeFMR08}; see also \\citealt{PlavchanBilinski13} for empirical evidence).  As tidal dissipation tends to damp the eccentricity while decreasing the orbital separation, close-in planets could then result from tidal circularization.  \\cite{WinnFAJ10} used a sample of 19 systems in which the projected spin-orbit angle $\\lambda$ was measured %with a $1\\sigma$ precision of $10^{o}$ or better via the Rossiter-McLaughlin (RM) effect. The RM effect occurs when a transiting planet blocks the blue- or red-shifted part (or both, depending on the orbital inclination with respect to the stellar spin) of the spinning star as it passes across the stellar disk, thus distorting the star's spectral line profile. \\cite{WinnFAJ10} investigated the behavior of the sky-projected misalignment as a function of the host star's effective temperature ($T_{\\rm eff *}$). Their results suggested that the degree of misalignment increases for hotter stars. In particular, a sharp increase in $\\lambda$ seems to occur at $T_{\\rm eff *}\\,\\simeq\\,$6250~K. These findings were later confirmed by \\citeauthor{Albrecht+12} (\\citeyear{Albrecht+12}, hereafter A12) with a sample of RM measurements twice as large as the one available to \\cite{WinnFAJ10}; see also e.g., \\cite{MortonJohnson11}. Since $\\simeq\\,$6250~K is the temperature at which the outer convective zone in main sequence stars starts becoming negligible, \\cite{WinnFAJ10} proposed that the mechanism responsible for the trend observed in the data is convective dissipation of tides in the star.  Hot Jupiters could then be produced via a single formation process yielding a broad distribution of obliquities. Later on, tidal dissipation in cool stars would damp the obliquity within a few Gyr. Correspondingly, the high degree of sky-projected misalignment observed in hot stars would result from the inefficiency of tidal dissipation. While its simplicity is appealing, this scenario presents a major weakness. In fact,  tidal dissipation in the star acts both on the misalignment and on the orbital separation, causing orbital decay whenever the stellar spin frequency is lower than the orbital frequency, and it thus fails in explaining the currently observed aligned hot Jupiters (see, e.g., \\citealt{RogersLin13}, hereafter R13). %\\cite{RogersLin13} (hereafter R13) integrated the equations describing the tidal evolution of the orbital separation, stellar spin and obliquity in the weak friction approximation considering both slowly and rapidly rotating stars. For slow rotators the authors find that, while tides tend to damp the obliquity, they also transfer angular momentum from the orbit to the star at an increasingly faster rate. As a result, the planet would fall into the star before the obliquity is completely damped. On the other hand, for fast rotators the authors show that tidal damping of the misalignment occurs together with a significant orbital expansion. A significant degree of obliquity evolution can occur only if the planet starts its orbital evolution from inside the star.  Possible solutions to this evolutionary conundrum were presented by \\cite{WinnFAJ10} and \\cite{Lai12}.  \\cite{WinnFAJ10}  suggested that, if the star's radiative interior is weakly coupled to the outer convective region and to the planet, then tides would act on the obliquity faster than on the orbital separation (see also A12). However, a large amount of differential rotation inside the star would potentially lead to fluid instabilities, which would tend to quickly re-couple the star's convective and radiative regions. To overcome this problem, \\cite{Lai12} presented a different scenario, following the idea that different physical processes dissipate tides with different efficiencies. His model invokes the excitation and damping of inertial waves in a stellar convective zone (see \\S~\\ref{Tidal Contribution due to inertial wave dissipation} for a summary). These waves are driven by the Coriolis force and are excited only in misaligned systems. In this configuration, the tidal potential to the leading quadrupole order has several terms. Each of these terms generates tidal disturbances with its own dissipation efficiency. Among these terms, \\cite{Lai12} identified a component of the tidal torque which acts only on the misalignment without affecting the orbital separation, thus providing a more efficient mechanism to modify the misalignment.  The validity of this prescription was recently questioned by R13, who  considered a random distribution of initial obliquities for 50 objects (nearly the number of observed systems considered by A12) and integrated the equations derived by \\cite{Lai12} forward in time. The authors computed the evolution of the misalignment {\\it alone} while keeping the orbit and stellar spin fixed and found that  tides would lead to a nearly equal amount of prograde ($\\lambda \\textless 90^{o}$),  retrograde ($\\lambda \\textgreater 90^{o}$), and 90$^{o}$ orbits. This appears inconsistent with the observations, as the majority of observed obliquities are smaller than $90^{o}$.  However, we note that R13's investigation has two major limitations. First, it neglects the simultaneous evolution of the orbital separation, stellar spin, and misalignment, while previous investigations have shown that it is essential to consider the coupled evolution of the orbital elements and spins (e.g., \\citealt{Jackson+08,BarkerOgilvie2009,MatsumuraPR2010,Xue+14}). Furthermore, it does not account for the various physical effects that might compete in the evolution of the system (e.g., magnetic braking and the radial expansion of the star as a result of stellar evolution). This last simplification was recently adopted also by \\cite{Xue+14}, who integrated the full set of equations presented by \\cite{Lai12} and showed that all intermediate states found by R13 eventually evolve towards alignment. While this appears to be consistent with the majority of observed obliquities ($\\textless\\,90^{o}$), it cannot explain the currently observed intermediate misalignments. %We show in this paper that the orbital configurations of misaligned hot Jupiters can indeed be explained, once all physical effects are properly taken into account.   In this paper, we reconsider tidal dissipation in the star as a possible mechanism responsible for the observed distribution of misalignments and orbital separations. We carefully examine the observed relation between $\\lambda$, $T_{\\rm eff *}$, and the amount of convection inside the host star. In contrast to previous studies, we use detailed stellar evolution models, thus accounting for the dependence of convection on stellar properties, such as mass, metallicity, effective temperature, and age. Furthermore, we integrate the full set of equations describing the evolution of the orbital separation, stellar spin, and misalignment. We take into account the effects of tidal dissipation in the star, stellar wind mass loss, changes in the star's internal structure as a result of stellar evolution, and magnetic braking. The tidal prescription adopted follows \\cite{Lai12}, and includes both tides in the weak friction approximation \\citep{Zahn1977, Zahn1989} and convective damping of inertial waves. In the weak-friction regime, a body's response to tides is generally measured via a tidal quality factor $Q$ \\citep{GoldreichSoter66}, which parametrizes the efficiency of tidal dissipation. This term measures how a tidally-deformed body undergoing a forced oscillation dissipates part of the associated energy during each oscillation period. It is formally defined as the ratio of the maximum energy stored in the tidal distortion over the energy lost during each cycle. %how a tidally deformed body undergoing a forced oscillation dissipates part of the associated energy during each oscillation period (similarly to the quality factor of a forced, damped harmonic oscillator). %, it is defined as %\\begin{align} %Q = 2\\pi E_{0}\\left(\\oint -\\dot{E}dt\\right)^{-1}. %\\end{align} %Here $E_{0}$ is the maximum energy stored in an oscillation and the integral is the energy dissipated over one cycle. The value of $Q$ is the result of complex dissipative processes occurring within a body and it thus varies for bodies of different masses and types. Furthermore, $Q$ depends on the tidal forcing frequency and thus on the spins and orbital configuration. As a result, $Q$ is expected to vary by orders of magnitudes \\citep{PenevSasselov11} and it is clear that different $Q$ values are needed to explain different systems (e.g., \\citealt{MatsumuraPR2010}, hereafter M10). In this work, we prefer not to introduce additional model parameters and instead use a parametrization for tidal dissipation calibrated from observations of binary stars (e.g.,  \\citealt{VerbuntPhinney1995,RasioTLL1996,HurleyTP02,BelczynskiStartrack2008}).  For quick reference, the notations adopted in this work for the components and orbital parameters are summarized in Table~\\ref{Tab:ParamsDefinition}.  The paper is organized as follows. In \\S~\\ref{The Sample} we present the sample of hot Jupiters considered in this work. In \\S~\\ref{Detailed Stellar modeling with MESA} we present the procedure adopted to model in detail the host stars in our sample. In \\S~\\ref{Orbital Evolution Model}, we summarize the equations that we integrate to study the orbital evolution of misaligned hot Jupiters (tests on the orbital evolution code developed for this work are presented in Appendix~\\ref{Tests on the Orbital Evolution code}). In \\S~\\ref{Detailed orbital evolution for four representative systems} we present possible evolutionary sequences of five representative systems: \\mbox{HAT-P-6}, WASP-7, 15, 16, and 71 (the results are summarized in Table~\\ref{Tab:OrbitalEvolResultsPandTheta}, where we also include a few additional examples, without describing their evolution in detail). We discuss the assumptions adopted in this work in \\S~\\ref{Discussion}. We summarize and conclude in \\S~\\ref{Conclusions}. As mentioned above, we consider tides in the weak friction approximation. Specifically, we account for both convective damping of the equilibrium tide and radiative damping of the dynamical tide. For the latter, we use results of detailed calculations presented by \\cite{Zahn1975}, which are valid in the limit of small tidal forcing frequencies (in the weak-friction regime). In Appendix~\\ref{Dynamic Tides in WASP-71} we solve the full set of equations describing non-adiabatic non-radial forced stellar oscillations, and we discuss the significance of dynamic tides in the most massive system among those studied in detail in \\S~\\ref{Detailed orbital evolution for four representative systems} for a wide spectrum of tidal forcing frequencies. %%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{table}[!h] \\centering \\caption{Definition of the various parameters used in this work.} \\begin{tabular}{lr} \\hline {\\bf Parameter} & {\\bf Definition} \\\\ \\\\[-1.0em] \\hline \\\\[-1.0em] \\smallskip $M_{*}$, $M_{\\rm pl}$, $\\Delta\\,M_{\\rm CZ}$ & Mass\\\\ \\smallskip $R_{*}$, $\\Delta\\,R_{\\rm CZ}$ & Radius\\\\ \\smallskip $L_{*}$ & Bolometric luminosity\\\\ \\smallskip $T_{\\rm eff *}$ & Effective temperature\\\\ \\smallskip Fe/H (or $Z$) &  Metallicity\\\\ \\smallskip $I_{*}$ &  Moment of inertia\\\\ \\smallskip $\\Omega_{*}$ ($\\Omega_{o}$) & Spin (orbital) frequency\\\\ \\smallskip $i_{*}$ ($i_{o}$) &   Stellar (orbital) inclination\\\\ \\smallskip $v_{{\\rm rot}}{\\rm sin}~i_{*}$ & Rotational velocity\\\\ \\smallskip $a$ & Semimajor axis\\\\ \\smallskip $P_{\\rm orb}$ & Orbital period\\\\ \\smallskip $e$ & Eccentricity\\\\ \\smallskip $S$ & Spin angular momentum\\\\ \\smallskip $L$ & Orbital angular momentum\\\\ $\\lambda$ ($\\Theta_{*}$) & Sky-projected (true) misalignment\\\\ \\\\[-1.0em] \\hline \\\\[-1.0em] \\end{tabular} \\label{Tab:ParamsDefinition} \\tablecomments{The subscripts ``*'' and ``pl'' refer to the star and planet, respectively, while ``CZ'' refers to the stellar surface convection zone (see \\S~\\ref{Detailed Stellar modeling with MESA} for details) . We denote with $i_{*}$  the angle between the stellar spin axis and the line of sight, while $i_{o}$ denotes the angle between the orbital angular momentum and the line of sight.} \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%% %In what follows we use $M_{*}$, $R_{*}$, $T_{\\rm eff *}$, Fe/H (or $Z$), $\\lambda$ ($\\Theta_{*}$), $I_{*}$, and $v_{{\\rm rot}}{\\rm sin}~i$ to indicate the star's mass, radius, effective temperature, metallicity, sky-projected (true) misalignment, moment of inertia, and rotational velocity, respectively. The planet's mass is denoted with $M_{\\rm pl}$. The stellar spin frequency and planet's orbital frequency are $\\Omega_{*}$ and $\\Omega_{o}$, respectively. The orbital period, semi major axis, and eccentricity are denoted with $P_{\\rm orb}$, $a$ and $e$, respectively. The angle between the stellar spin axis and the line of sight is $i_{*}$, while the angle between the orbital angular momentum and the line of sight is $i_{o}$. %%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "Two formation models have been proposed to explain the tight orbits of hot Jupiters. These giant planets could migrate inward in a disk (in the so-called {\\it disk migration} scenario), or they could be formed via tidal circularization of a highly eccentric orbit following gravitational interactions with a companion (in the so-called {\\it high-eccentricity migration} scenario). Disk migration yields orbits where the stellar spin and orbital angular momentum vectors are nearly aligned, while high-eccentricity migration results in a broad range of misalignments. Here we targeted the known hot Jupiters where the obliquity has been inferred observationally (following and updating the sample considered by \\citealt{Albrecht+12}) and investigated whether their properties are consistent with high-eccentricity migration.  In contrast to previous studies, we modeled in detail each host star and showed that the observed increase in misalignment $\\lambda$ with the star's effective temperature $T_{\\rm eff}$ is shaped by the amount of convection inside the host star, as originally suggested by \\cite{WinnFAJ10}. Specifically, higher degrees of misalignment are found in stars with less surface convection, especially if one considers the radial extent of the surface convective region. This result supports the hypothesis that giant planets are formed with a broad initial distribution of misalignments. During the subsequent orbital evolution, convective dissipation of tides in the star is the mechanism that shapes the observed distribution of misalignments. To further test this hypothesis, we computed the coupled evolution of the orbital elements and stellar spin of five representative systems: one aligned, two prograde, and two retrograde systems. We studied in detail \\mbox{HAT-P-6}, WASP-7, 15, 16, and 71, and provided results for few more systems, as additional examples. Within the regime of validity of the tidal prescription adopted (\\S~\\ref{Tidal Contribution due to inertial wave dissipation}), WASP-16 is an aligned star among the coolest ones. Instead, \\mbox{WASP-71} and \\mbox{WASP-7} (\\mbox{WASP-15} and \\mbox{HAT-P-6}) are among the coolest and hottest prograde (retrograde) stellar hosts, respectively. In contrast to previous studies on the observed  $\\lambda\\,-\\,T_{\\rm eff *}$ correlation (e.g., \\citealt{Albrecht+12,RogersLin13}), we took into account the combined effects of tides, stellar wind mass loss, magnetic braking, and stellar evolution. The tidal prescription adopted combines tides in the weak friction approximation and inertial wave dissipation, and it was recently proposed by \\cite{Lai12} to explain the currently observed aligned hot Jupiters. For the efficiency of inertial wave dissipation, we followed numerical results by \\cite{OgilvieLin2007} and \\cite{BarkerOgilvie2009} and considered tidal quality factors $Q'_{\\rm 10}$ ranging between $10^{6}\\,-\\,10^{10}$. Furthermore, we scanned a broad parameter space in initial orbital periods, misalignments, and degrees of asynchronism between the stellar spin and the planet's orbital frequency ($\\Omega_{*}/\\Omega_{o}$).  Our results show that, accounting for all the relevant physical mechanisms and considering the simultaneous evolution of the orbital separation, stellar spin, and misalignment, the current properties of the variety of systems considered here can be naturally explained.  For HAT-P-6, WASP-7, and 15, with F-dwarfs, we found that both orbital decay and obliquity damping are small even on Gyr timescales. This supports the notion that many of the F stars are presently not capable of either destroying the planet or damping the obliquity, and that the high observed obliquities are the result of the hot-Jupiters formation process.  For WASP-71, we found that both the orbit and obliquity are actively damping, but, for the smallest $Q'_{\\rm 10}$ value considered, the obliquity evolves on a timescale that can be almost 1 order of magnitude shorter at present (and up to 3 orders of magnitude shorter in the past, Fig.~\\ref{fig:DetailedOrbitalEv_WASP-71_timescales}). The same is true for WASP-16, where the obliquity damping timescale is more than 1 order of magnitude shorter than the orbital decay timescale at present (Fig.~\\ref{fig:DetailedOrbitalEv_WASP-16_timescales}). This system in particular, together with WASP-4 described in \\cite{VR14}, supports the idea that obliquity damping in cool and less massive G-dwarfs can occur more rapidly than orbital decay, provided that inertial wave dissipation is actively driving the misalignment evolution. This mechanism does indeed provide an explanation for the population of currently known aligned hot Jupiters.  Finally, our results show that the physical effects included in this work can all play a significant role in the orbital evolution of misaligned hot Jupiters systems, depending on the properties of the star and planet, as well as the orbital configuration, as summarized below.  By using detailed stellar models we took into account how changes in the star's moment of inertia affect its spin. Stellar evolution efficiently decreases $\\Omega_{*}$ in systems hosting F-dwarfs, where it becomes increasingly significant as the star expands during its main sequence evolution. Another efficient driver of stellar spin-down is magnetic braking, included here according to \\citeauthor{Skumanich72}'s (\\citeyear{Skumanich72}) law. We varied its strength by changing the parameter $\\gamma_{\\rm MB}$ entering the magnetic braking prescription, but the values of $\\gamma_{\\rm MB}$ which yield \\mbox{HAT-P-6-}, \\mbox{WASP-15-}, and \\mbox{71-}type systems agree with those adopted in the literature. WASP-7, and 16 are exceptions, as their present properties can be matched either {\\it without} magnetic braking ($\\gamma_{\\rm MB}\\,=\\,0$, for WASP-7, with an F-dwarf) or with $\\gamma_{\\rm MB}\\,\\leq\\,0.5$ (for WASP-16, with a G-dwarf). However, we argued that more systems resembling WASP-7 with $\\gamma_{\\rm MB}\\,>\\,0$ could be found considering higher initial levels of asynchronism between the stellar spin and the planet's orbital frequency.  For tides, the \\citeauthor{Lai12}'s (\\citeyear{Lai12}) prescription is valid for sub-synchronous systems ($\\Omega_{*}\\,\\textless\\,\\Omega_{o}$) and inertial wave dissipation acts only on the stellar spin and the misalignment. Instead, in super-synchronous systems ($\\Omega_{*}\\,\\textgreater\\,\\Omega_{o}$), this dissipation channel might also affect the orbital separation, but we did not explore this possibility and we targeted only configurations where $\\Omega_{*}\\,\\leq\\,\\Omega_{o}$ (and discussed cases in which $\\Omega_{*}\\,=\\,\\Omega_{o}$). %The results do not change significantly for $Q'_{\\rm 10}\\textgreater 10^{6}$, for most system configurations. %In fact, explaining the currently observed misaligned systems {\\it solely} via tides in the weak-friction approximation presents major problems (e.g., \\citealt{RogersLin13}). %This tidal prescription was recently questioned by R13, however, their investigation present major weaknesses. In fact, while it is clear from the literature (e.g., this paper, \\citealt{Jackson+08,BarkerOgilvie2009,MatsumuraPR2010}) that it is essential to consider the coupled evolution of the orbital and rotational elements, R12 computed the evolution of the misalignment {\\it alone, only} accounting for inertial wave dissipation, thus neglecting all the various and potentially significant physical effects included here. As mentioned above, we showed that inertial wave dissipation can be significant, mainly for the evolution of the misalignment. This dissipation channel can increase the misalignment in systems like \\mbox{HAT-P-6}, while it can actively damp it in WASP-16- and 71-type systems. It is unimportant for WASP-15- and 7-type systems and the associated timescales can be orders of magnitude longer than those related to the main drivers of spin and misalignment evolution. For systems like \\mbox{WASP-7}, in particular, this finding suggests that the evolution of the orbital separation would not be significantly affected by inertial wave dissipation, if super-synchronous configurations were considered. Expanding the initial parameter space to include $\\Omega_{*}/\\Omega_{o}\\,\\textgreater\\,1$ would likely yield \\mbox{WASP-7-type} systems where the magnetic braking coefficient $\\gamma_{\\rm MB}$ is consistent with values adopted in the literature.  Tides in the weak friction approximation are the main driver of orbital decay (when $\\Omega_{*}\\,\\textless\\,\\Omega_{o}$) and they can also affect the misalignment. As weak-friction tides always tend to decrease the misalignment, they can either counteract the effect of inertial wave dissipation (e.g., in \\mbox{HAT-P-6-type} systems) or strengthen it (e.g., in \\mbox{WASP-71-type} systems), depending on $Q'_{\\rm 10}$. Tides in the weak friction approximation are less important for the evolution of the stellar spin, except for \\mbox{WASP-71}: for this system, this mechanism becomes relevant towards the end of the evolution and it counteracts the spin-down driven by stellar evolution and magnetic braking.  Finally, stellar wind mass loss is negligible for all systems considered here. In a \\mbox{WASP-71-type} system it could have been the main mechanism driving the evolution of the orbital separation early in its past, but the associated timescales were too long to significantly affect it.  To conclude, our detailed stellar modeling, the orbital evolution calculations, and the few more examples summarized in Table~\\ref{Tab:OrbitalEvolResultsPandTheta}, provide support to the high-eccentricity migration scenario for the formation of hot Jupiters. A detailed calculation similar to the one presented here on {\\it all} misaligned systems is needed to provide a definite answer (this will be the subject of future work, see also \\S~\\ref{Discussion}). However, here we demonstrated that tidal dissipation in the host star, together with stellar wind mass loss, magnetic braking, and stellar evolution, can account for the observed distribution of misalignments (and alignments) and orbital separations found around stars of different temperatures. In particular, the orbital configuration of five representative hot Jupiter systems covering a variety of misalignments and stellar temperatures can be explained with reasonable assumptions on the physical effects driving the orbital evolution in circular and misaligned systems, despite the limited parameter space considered."
    },
    "1402/1402.2342_arXiv.txt": {
        "abstract": "We report resolved near-infrared spectroscopic monitoring of the nearby L dwarf/T dwarf binary WISE~J104915.57$-$531906.1AB (Luhman~16AB), as part of a broader campaign to characterize the spectral energy distribution and temporal variability of this system. A continuous 45-minute sequence of low-resolution IRTF/SpeX data spanning  0.8--2.4~$\\micron$ were obtained, concurrent with combined-light optical photometry with ESO/TRAPPIST. Our spectral observations confirm the flux reversal of this binary, and we detect a wavelength-dependent decline in the relative spectral fluxes of the two components coincident with a decline in the combined-light optical brightness of the system over the course of the observation. These data { are successfully modeled as a combination of achromatic (brightness) and chromatic (color) variability in the T0.5 Luhman~16B, consistent with variations in overall cloud opacity; and no significant variability in L7.5 Luhman~16A, consistent with recent resolved photometric monitoring.} We estimate a peak-to-peak amplitude of 13.5\\% at 1.25~$\\micron$ over the full lightcurve. Using a simple { two-spot} brightness temperature model for Luhman~16B, we infer { an average cold} covering fraction of { $\\approx$30--55\\%, varying by 15--30\\% over a rotation period assuming a $\\approx$200--400~K difference between hot and cold regions. We interpret these variations as changes in the covering fraction of a high cloud deck and corresponding ``holes''  which expose deeper, hotter cloud layers, although other physical interpretations are possible. A Rhines scale interpretation for the size of the variable features explains an apparent correlation between period and amplitude for Luhman~16B and the variable T dwarfs SIMP 0136+0933 and 2MASS~J2139+0220, and predicts relatively fast winds (1--3~{\\kms}) for Luhman~16B consistent with lightcurve evolution on an advective time scale (1--3 rotation periods). The strong variability observed in this flux reversal brown dwarf pair supports the model of a patchy disruption of the mineral cloud layer as a universal feature of the L dwarf/T dwarf transition.} ",
        "introduction": "The driving mechanism for the transition between the L dwarf and T dwarf spectral classes has emerged as one of the outstanding problems in brown dwarf astrophysics. Spectroscopically, this transition is defined by the { appearance} of {\\meth} absorption features at near-infrared wavelengths \\citep{1999ApJ...519..802K,2006ApJ...637.1067B}, accompanied by a substantial reduction of condensate cloud opacity \\citep{1996Sci...272.1919M,2000ApJ...531..438B,2001ApJ...556..357A}. Both effects drive near-infrared spectral energy distributions (SEDs) to transition from red ($J-K \\approx 1.5-2.5$) to blue ($J-K \\approx 0-0.5$; \\citealt{2000ApJ...536L..35L,2002ApJ...568..335M,2006ApJ...637.1067B}), with strengthening molecular gas bands delineating the spectral subclasses. What is remarkable about the L dwarf/T dwarf transition is that it appears to take place over a relatively narrow range of effective temperatures ({\\teff}s) and luminosities, based on absolute magnitude trends (e.g., \\citealt{2002AJ....124.1170D,2004AJ....127.2948V}), broad-band SED measurements (e.g., \\citealt{2004AJ....127.3516G}) and spectral model fits (e.g., \\citealt{2008ApJ...678.1372C,2009ApJ...702..154S}). The L/T transition also exhibits an apparent excess of binaries \\citep{2007ApJ...659..655B}, gaps in color distributions \\citep{2012ApJS..201...19D} and a decline in number densities as a function of spectral type \\citep{2008ApJ...676.1281M}, trends that suggest the transition is rapid in time as well as temperature.  The important role of photospheric cloud evolution for this transition is seen in the observation that early-type T dwarfs with minimal cloud opacity are often significantly brighter at 1~$\\micron$ than their hotter, cloudier L dwarf counterparts. This is true in both color-magnitude diagrams of local populations \\citep{2003AJ....126..975T,2012ApJ...752...56F,2012ApJS..201...19D} and among components of ``flux-reversal'' binaries that straddle the L/T transition \\citep{2006ApJS..166..585B,2006ApJ...647.1393L,2008ApJ...685.1183L}. The 1~$\\micron$ region is a minimum of molecular gas opacity---the local pseudocontinuum---so condensate grain scattering can dominate the overall opacity at these wavelengths \\citep{2001ApJ...556..872A}. The 1~$\\micron$ brightening has thus been interpreted as a depletion of photospheric condensate clouds over a narrow range of {\\teff} and/or time. The geometry of the depletion has been modeled as both global changes in photospheric chemistry (e.g., \\citealt{2004AJ....127.3553K,2005ApJ...621.1033T,2006ApJ...640.1063B,2008ApJ...689.1327S}) and hole formation that allows light to emerge from hotter regions \\citep{2001ApJ...556..872A,2002ApJ...571L.151B,2010ApJ...723L.117M}. The latter hypothesis predicts an enhancement of rotationally-modulated photometric variability at the L/T transition, particularly in the 1~$\\micron$ region, { depending on the sizes and distribution of the cloud gaps}.  Recent brown dwarf monitoring observations support this prediction, as the two most prominent variables identified to date, SIMP~J013656.5+093347 (\\citealt{2006ApJ...651L..57A,2009ApJ...701.1534A}; hereafter SIMP~J0136+0933) and 2MASS J21392676+0220226 (\\citealt{2008AJ....136.1290R,2012ApJ...750..105R}; hereafter 2MASS~J2139+0220) are both early-type T dwarfs. Their variability can be reproduced with spot models assuming regions with thick and thin clouds at different temperatures { assumed to probe different layers in the atmosphere}   \\citep{2012ApJ...750..105R,2013ApJ...768..121A}.  The spectral character of the observed variability is nevertheless complex.  Rather than variability being limited to pseudocontinuum regions where gas opacity is a minimum, broad-band chromatic and achromatic variations are seen across the infrared \\citep{2009ApJ...701.1534A,2012ApJ...750..105R}. The light curve shapes themselves are also seen to change over several rotation periods, suggesting dynamic evolution of features at rates considerably faster than the Solar gas giants \\citep{2013ApJ...776...85S}.  Finally, variability measurements over widely-separated spectral regions have recently revealed evidence of pressure-dependent phase variations, indicating vertical structure in { the features driving the variability} \\citep{2012ApJ...760L..31B,2013ApJ...778L..10B}.  The considerable level of detail on brown dwarf cloud structure and atmospheric dynamics garnered from these monitoring studies is of relevance to exoplanet atmospheres, where clouds are { now seen} as a key opacity source \\citep{2011ApJ...733...65B,2011ApJ...737...34M,2012ApJ...754..135M,2013MNRAS.432.2917P,2013A&A...559A..33C}.   The recently-discovered, nearby binary brown dwarf system {\\name}AB (hereafter {\\namesh}AB; \\citealt{2013ApJ...767L...1L}) has emerged as a potential benchmark for studying the L/T transition. With spectral types of L7.5 and T0.5 \\citep{2013ApJ...770..124K,2013ApJ...772..129B}, its components straddle the transition. Its T dwarf secondary is brighter than the primary in the 0.95--1.3~$\\micron$ range, making it a flux-reversal system \\citep{2013ApJ...772..129B}.  {\\namesh}AB is also a significant variable. Combined-light red optical photometry by \\citet{2013A&A...555L...5G} revealed peak-to-peak variability of $\\sim$10\\% with a period of 4.87$\\pm$0.01~hr, with large changes in the light curve structure over daily timescales.  The variability was attributed primarily to the T dwarf component. Resolved photometry by \\citet{2013ApJ...778L..10B} extended the observed variability into the near-infrared, confirmed {\\namesh}B as the dominant variable, and revealed pressure-dependent phase variations.  As such, this system embodies nearly all of the remarkable characteristics of the L/T transition---multiplicity, variability, and flux reversal---while residing only 2.020$\\pm$0.019~pc from the Sun \\citep{2013arXiv1312.1303B}.  In April 2013, our consortium organized a week-long monitoring campaign of {\\namesh}AB using telescopes in Chile, Australia and Hawaii, with the aim of characterizing its variability panchromatically (radio, optical and infrared) and spectroscopically, while simultaneously obtaining kinematic data (radial and rotational velocities) to constrain its orbit and viewing geometry.  This article reports low-resolution near-infrared spectroscopic monitoring observations obtained over 45 minutes with the SpeX spectrometer \\citep{2003PASP..115..362R} on the 3.0m NASA Infrared Telescope Facility (IRTF), coincident with combined-light optical photometry obtained with the TRAnsiting Planets and PlanetesImals Small Telescope (TRAPPIST; \\citealt{2011Msngr.145....2J}). In Section~2 we describe our observation and data reduction procedures, including period analysis of the TRAPPIST lightcurve around this epoch. In Section~3 we describe our spectral extraction and variability analysis of the SpeX data, and create an empirical model to replicate both the SpeX and TRAPPIST observations. In Section~4 we discuss our results, examining the nature of {\\namesh}B's inferred variability in the context of a simple { two-spot} brightness temperature model, and compare this source to other significantly variable L/T transition objects. We summarize our results in Section~5. ",
        "conclusions": "\\subsection{{ Interpreting} the Nature of {\\namesh}B's Spectral Variability}  Spectral trends in variability have been examined in several L and T dwarfs to date, through pure spectroscopy (e.g., \\citealt{2008MNRAS.384.1145B,2008A&A...487..277G,2013ApJ...768..121A}) and simultaneous or near-simultaneous broad-band imaging (e.g., \\citealt{2004MNRAS.354..466K,2009ApJ...701.1534A,2012ApJ...760L..31B,2012ApJ...750..105R,2013AJ....145...71K,2013ApJ...767..173H,2013ApJ...778L..10B}).  The most significant variables up until now, SIMP~J0136+0933 and 2MASS~J2139$-$0220, both exhibit color trends in { near-infrared} photometric variability, with larger amplitude changes at $J$ as compared to $K$, again consistent with variable condensate cloud opacity.  However, spectroscopic variability measurements of these same two sources over 1.0--1.7~$\\micron$ by \\citet{2013ApJ...768..121A} indicate that achromatic or near-achromatic variations dominate the psuedocontinuum.  These authors propose { a two-layer cloud model with a thick shallow cloud and thin deep cloud as a means of reproducing} both achromatic psuedocontinuum and chromatic broadband variability. { Matched to atmosphere models, this framework can replicate observed trends in the colors and spectral shapes of SIMP~J0136+0933 and 2MASS~J2139$-$0220 over 1-3 rotation periods, although detailed fits to the data remain poor (see also \\citealt{2012ApJ...750..105R}).}  For {\\namesh}B, we also find that both achromatic and chromatic variations must be present in the psuedocontinuum to properly model the observations.  Achromatic variation yields a decline in the overall flux, amounting to roughly 0.03~mag in broad-band $J$ over the observing period. The concurrent chromatic variation simultaneously reddens the spectrum of this source by $\\Delta(J-K)$ = 0.02~mag, resulting in a relative flux variation amplitude of $\\Delta{F_{K_s}}/\\Delta{F_J}$ = 0.41$\\pm$0.18, similar to values reported for SIMP~J0136+0933 \\citep{2009ApJ...701.1534A} and 2MASS~J2139$-$0220 \\citep{2012ApJ...750..105R}. Combined, the achromatic and chromatic terms nearly cancel in the $K$-band, a region that is gas opacity dominated (H$_2$O, CH$_4$ and H$_2$). Hence, { our linear spectral model} is functionally consistent with condensate clouds being the primary driver of variability in {\\namesh}B.  \\subsection{ A Brightness Temperature Spot Model for Luhman~16B}  { Given the known shortcomings in reproducing the near-infrared spectra of L/T transition brown dwarfs (e.g., \\citealt{2008ApJ...682.1256L,2008ApJ...678.1372C,2009ApJ...702..154S}), we forgo detailed modeling of the spectra in lieu of a simply brightness temperature variation model, focusing at 1.25~$\\micron$ where gas opacity is a minimum and cloud structure variations are expected to have the greatest influence \\citep{2001ApJ...556..872A}. The simplest model for replicating the surface flux $\\langle{F}\\rangle$ of a patchy brown dwarf is two sets of regions with differing brightness temperatures covering the surface:} \\begin{equation} \\langle{F}\\rangle \\propto \\langle{T_{br}}^4\\rangle \\equiv {A}T_{cold}^4 + (1-A)T_{hot}^4. \\label{eqn:tsurf} \\end{equation} { Here, $\\langle{T_{br}}^4\\rangle$ is the disk-averaged brightness temperature,} $T_{cold}$ and $T_{hot}$ are the brightness temperatures of cold and hot { regions}, respectively, $A \\equiv F_{cold}/\\langle{F}\\rangle \\leq 1$ is the areal covering fraction of the cold { regions}, and we ignore limb darkening. Our interpretation of this model is that the cold { regions} correspond to the { highest cloud layer in the brown dwarf atmosphere}, while the hot { regions} correspond to gaps in these clouds  that probe to some as-yet undetermined deeper layer with brightness temperature $T_{hot}$; { a cartoon perspective of this is shown in Figure 6 of \\citet{2013ApJ...768..121A}. We note that this is not the only interpretation of a two-spot model, which could also arise from magnetic interaction at the photosphere (i.e., starspots) or updrafts of warm air pockets driven by convective flows. Nevertheless, we will occasionally refer to the cold { region} as ``clouds'' and hot { regions} as ``holes'' in the following discussion.}  \\begin{figure}[h] \\center \\epsscale{1.0} \\plotone{f5.eps} \\caption{{ Cold spot} covering fraction $A$ and change in covering fraction $\\Delta{A}$ (solid lines) on {\\namesh}B { over a full cycle}, as a function of the hot spot brightness temperature at 1.25~$\\micron$. Values for $A$ are derived from Eqn.~\\ref{eqn:tsurf} and assume $\\langle{T_{br}}\\rangle_{L16B}$ = { 1560~K} and $T_{cold}$ = ${\\langle}T_{br}\\rangle_{L16A}$ = { 1510~K}, based on Figure~\\ref{fig:variability} (vertical triple-dot dash lines).  Values for $\\Delta{A}$ are derived from  Eqns.~\\ref{eqn:da1} and~\\ref{eqn:da2} and assume $\\Delta{F}/F$ = 13.5\\%.  The point where $\\Delta{A}$ = 100\\% indicates the minimum $T_{hot}$; { the temperature at which $A = 0.5$ is also labeled, separating cold surfaces with hot spots from a hot surface with cold spots.} Also labeled are the range of $T_{cr}$ values from \\citet{2005ApJ...621.1033T} and \\citet{2012ApJ...760..151S} that estimate the height of cloud tops, and the temperature at which all condensates are assumed to be evaporated ($T_{evap}$).  Estimates for $A$ and $T_{hot}$ for SIMP~J0136+0933 and 2MASS~J2139$-$0220 from \\citet{2012ApJ...750..105R} and \\citet{2013ApJ...768..121A} are indicated, assuming the same cloud top temperature.  Finally, we label estimates of jet size scales for wind velocities of { $U$ = 1.6~{\\kms} and 3.4~{\\kms} based on a Rhines length scale (Eqn.~\\ref{eqn:rhines}); these intersect the $\\Delta{A}$ curve at $T_{hot}$ = 1700~K and 1900~K}. } \\label{fig:patchy} \\end{figure}   For {\\namesh}B,  ${\\langle}T_{br}\\rangle$ = { 1560~K} at the 1.25~$\\micron$ $J$-band peak continuum (Figure~\\ref{fig:variability}; see also \\citealt{faherty14}).  If we take the brightness temperature of {\\namesh}A at this wavelength, { 1510}~K, as an estimate for $T_{cold}$ for both sources,\\footnote{The assumption can be justified in part by the nearly identical brightness temperatures of the two sources in the 1.15~$\\micron$ region, where {\\wat} and {\\meth} opacity play a larger role than clouds.} then we can jointly constrain $A$ and $T_{hot}$, as illustrated in Figure~\\ref{fig:patchy}.  { Coverage of cold regions} is essentially negligible for $T_{hot} <$ 1570~K, then climbs to over 50\\% at $T_{hot} \\approx$ 1860~K.  At hotter temperatures, { our model predicts that the atmosphere of {\\namesh}B would be overall similar to that of {\\namesh}A with occasional hot spots, which we assume to be less than the evaporation temperature of mineral condensate species} ($T_{evap} \\approx$ 2000~K; \\citealt{1999ApJ...519..793L}). We note that { equal hot-cold spot coverage for {\\namesh}B  occurs in} the 1700--1900~K range that \\citet{2005ApJ...621.1033T} estimate as the effective top of a brown dwarf cloud layer ($T_{cr}$; see \\citealt{2012ApJ...760..151S}). \\citet{2012ApJ...750..105R} and \\citet{2013ApJ...768..121A} also provide estimates for $A$ and $\\Delta{T}_{hc} = T_{hot}-T_{cold}$ for SIMP~J0136+0933 and 2MASS~J2139$-$0220 based on their own { two-spot modeling}. While temperature values reported in these studies are based on model effective temperatures, if we assume that the brightness temperature offsets are the same as their preferred $\\Delta{T}_{hc} \\approx$ 300~K, this would also place the hot { regions of {\\namesh}B in the same temperature range as} the cloud tops of the \\citet{2005ApJ...621.1033T} models. { Thus, if our spot model is interpreted as probing different layers of {\\namesh}B's atmosphere, the best estimates of the temperature differential is in line with the conjecture of \\citet{2013ApJ...768..121A} that gaps in the highest cloud deck still probe regions influenced by condensate opacity. However, we stress that our data cannot independently determine $A$ or $T_{hot}$, and other interpretations of these temperature differences are conceivable.}   The fractional peak-to-peak variation in observed flux that occurs as { hot and cold regions} rotate in and out of view is \\begin{equation} \\frac{\\Delta{F}}{\\langle{F}\\rangle} = \\frac{\\Delta{A}}{A-\\epsilon} \\label{eqn:da1} \\end{equation} where $\\langle{F}\\rangle$ is the average flux, $\\Delta{F}$ the change in total flux, $\\Delta{A}$ the change in cloud coverage (increasing $A$ decreases the total flux), and $\\epsilon \\equiv F_{hot}/(F_{hot}-F_{cold}) = T_{hot}^4/(T_{hot}^4-T_{cold}^4) \\geq 1$; see also \\citet{2012ApJ...750..105R}.\\footnote{Our expression differs slightly from \\citet{2012ApJ...750..105R} because we assign $A$ to be the cloud-covering fraction, whereas they define the equivalent parameter $a$ as the cloud-cleared fraction.} In terms of brightness temperatures: \\begin{equation} \\Delta{T_{br}}  =  \\frac{\\Delta{F}}{4\\langle{F}\\rangle}\\langle{T_{br}}\\rangle = \\frac{\\Delta{A}}{4(A-\\epsilon)}\\langle{T_{br}}\\rangle. \\label{eqn:da2} \\end{equation} During our observations, we observed a 7.5\\% variation in the $J$-band peak continuum that was coincident with a 2.5\\% variation in TRAPPIST red-optical photometry. We therefore assume that the full 4.5\\% peak-to-peak variation in TRAPPIST photometry around our spectral observations (Figure~\\ref{fig:trappist}) corresponds to a 13.5\\% variation at $J$, or a peak-to-peak temperature fluctuation of $\\Delta{T_{br}} \\approx$ 50~K\\footnote{The maximum variation observed by \\citet{2013A&A...555L...5G} over weeks of monitoring is 10\\%, which would correspond to 30\\% variations in $J$, exceeding those observed in 2MASS~J2139$-$0220. However, we restrict our analysis here to the period around the spectral observations { since the spectral response of larger fluctuations may differ}.} following Eqn~\\ref{eqn:da2}.  This temperature offset is notably similar to the temperature difference between {\\namesh}A and B at these wavelengths (Figure~\\ref{fig:variability}). Using the { relationship between} $A$ and $T_{hot}$ above, we computed $\\Delta{A}$ as a function of $T_{hot}$, also shown in Figure~\\ref{fig:patchy}.  Not surprisingly, the { areal} variation required to reproduce the observed brightness variations declines with higher $T_{hot}$; i.e., with greater contrast between cold and hot { regions}. An important reference point is the temperature at which { areal} variations become smaller than the total { cold region} coverage, which occurs for $T_{hot} >$ 1710~K and $A >$ 30\\%. { The corresponding $\\Delta{T}_{hc}$ = 150~K is on the low end of estimates} for  SIMP~J0136+0933 and 2MASS~J2139$-$0220, and just above the minimum $T_{cr}$ from \\citet{2005ApJ...621.1033T}. { For the range 1700~K $< T_{hot} <$ 1900~K, which we again take as a reasonable estimates of the hot spot temperature, the inferred cold covering fraction is roughly 30--55\\%, intermediate between similar values inferred for SIMP~J0136+0933 (25-30\\%) and 2MASS~J2139$-$0220 (50-65\\%) by \\citet{2013ApJ...768..121A}.   \\subsection{ Interpretation: Rhines Length Scale and Advective Time Scale}  For 1700~K $< T_{hot} <$ 1900~K, { cold spot coverage} must vary by 15--30\\% over a single period to replicate the observed variability amplitude, { implying a $\\sim$ 30-100\\% variation between hemispheres if the spot patterns are static. Organized jet features in the atmospheres of the giant Solar planets generally scale in size with the Rhines length \\citep{1970GApFD...1..273R,2008ASPC..398..419S}, $L_{Rh} \\sim \\left({U}/{2\\Omega{R}\\cos{\\phi}}\\right)^{1/2}$, where $U$ the characteristic wind speed, $R$ is the radius, $\\Omega= 2\\pi/P$, $P$ is the rotation period and $\\phi$ is the latitude of the feature.  If we assume that the same scaling occurs for features in brown dwarf atmospheres (e.g., \\citealt{2013ApJ...768..121A,2013ApJ...776...85S}), then their maximum fractional size scale is: \\begin{equation} \\alpha_{Rh} \\sim \\left(\\frac{L_{Rh}}{R}\\right)^2 \\approx  2\\%\\left(\\frac{U}{\\rm km/s}\\right)\\left(\\frac{P}{\\rm hr}\\right)\\left(\\frac{R_{Jup}}{R}\\right). \\label{eqn:rhines} \\end{equation} where we have assumed mid-latitude features. If we now relate this maximum scale to the areal spot variation inferred here ($\\alpha_{Rh} \\sim \\Delta{A}$), the known rotational period and assumed radius of {\\namesh}B implies characteristic wind speeds of 1.6~{\\kms} $< U <$ 3.4~{\\kms} for 1700~K $< T_{hot} <$ 1900~K (Figure~\\ref{fig:patchy}). These speeds are somewhat higher than the range favored by the circulation models of \\citet{2013ApJ...776...85S}, assuming winds are driven by inefficient conversion of convective heat (10--300~{\\ms})}. However, the speeds do give advection timescales, $\\tau_{adv} \\sim R/U \\sim (2-5)\\times10^4~s \\sim 1-3$ rotation periods, that are consistent with the timescale of lightcurve evolution observed in {\\namesh}B  \\citep{2013A&A...555L...5G}.  The convergence between the inferred variation and Rhines length scales, and the advective and evolutionary time scales, suggest that our gross estimates for $T_{cold}$, $T_{hot}$, $A$ and $\\Delta{A}$ are not too far off the mark.  However, we have made a number of major assumptions that require confirmation through more detailed spectroscopic monitoring and modeling, in particular to ascertain whether the spot regions have spectral characteristics (features and line profile shapes) consistent with the inferred brightness temperatures.  Nonetheless, our basic model of a cold cloud deck disrupted by warm dynamic features shows promising agreement with planetary analogs and current brown dwarf circulation models.}  \\subsection{Trends in L/T Transition Variability}  \\begin{deluxetable*}{lcccccl} \\tablecaption{Comparison of Highly Variable L/T Transition Dwarfs. \\label{tab:comp}} \\tabletypesize{\\small} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{SpT} & \\colhead{$P_{rot}$} & \\colhead{$\\Delta{F}/F$} & \\colhead{$A$} & \\colhead{$\\Delta{K}/\\Delta{J}$} & \\colhead{Ref} \\\\ & & \\colhead{(hr)} & \\colhead{at 1.25~$\\micron$} \\\\ } \\startdata SIMP~J0136+0933\t& T2.5 & 2.3895$\\pm$0.0005 & 5.5\\% & 25-30\\% & 0.48$\\pm$0.06 & 1,2 \\\\ {\\namesh}B  & T0.5 & 4.87$\\pm$0.01 & 13.5\\%\\tablenotemark{a} &  30-55\\%\\tablenotemark{b} & 0.41$\\pm$0.18 &  3,4 \\\\ 2MASS~J2139$-$0220 & T1.5 & 7.721$\\pm$0.005 & 30\\% & 50-65\\% & 0.45--0.83 & 2,5 \\\\ \\enddata \\tablenotetext{a}{Based on the maximum peak-to-peak TRAPPIST variability amplitude during the current observing period.} \\tablenotetext{b}{Assuming $T_{hot}-T_{cold}$ = 300$\\pm$100~K; see \\citet{2013ApJ...768..121A}.} \\tablerefs{ (1) \\citet{2009ApJ...701.1534A}; (2) \\citet{2013ApJ...768..121A}; (3) \\citet{2013A&A...555L...5G}; (4) This paper; (5) \\citet{2012ApJ...750..105R}.} \\end{deluxetable*}   {\\namesh}B joins SIMP~J0136+0933 and 2MASS~J2139$-$0220 as the three most variable L/T transition objects detected to date, so it is worth comparing the variability properties of these sources, summarized in Table~\\ref{tab:comp}.   The variability period, $J$-band variability amplitude and inferred cloud covering fraction of {\\namesh}B are all intermediate between those of SIMP~J0136+0933 and 2MASS~J2139$-$0220, although epoch-to-epoch changes in these values are considerable.  As the Rhines scale scales linearly with the rotation period,\\footnote{\\citet{2013ApJ...768..121A} incorrectly state a spot scaling law of ${P}^{-2}$ in the text, but infer a spot scaling between 2MASS~J2139$-$0220 and SIMP~J0136+0933 that is consistent with $A \\propto P$; the former is likely a typographical error.}, { its interpretation as an estimate of surface feature size is consistent with {\\namesh}B's intermediate period and intermediate variability amplitude, as a few large features are more likely to give rise to stronger disk-integrated variations than many small features \\citep{2013ApJ...768..121A}.} There also appears to be a correlation between rotation period and cloud covering fraction, although temperature effects may play a role in this statistic. The source with the smallest cloud coverage, SIMP~J0136+0933, is also the latest-type and presumably coldest brown dwarf in the sample.  Finally, we find essentially no difference in color variability among these sources. As noted above, our estimate of $\\Delta{F_{K_s}}/\\Delta{F_{J}}$ for {\\namesh}B is consistent with similar measures for SIMP~J0136+0933 and 2MASS~J2139$-$0220 { (although the latter can exhibit more extreme color terms; \\citealt{2012ApJ...750..105R}), suggesting that the condensate clouds responsible for the variations in these sources are likely to} have similar opacities and physical properties (i.e., composition, grain size distribution, vertical structure, etc.). { However, confirmation of this agreement will again require more careful spectral modeling to accurately determine cloud properties.}"
    },
    "1402/1402.5081_arXiv.txt": {
        "abstract": "{% In Monte-Carlo simulations of gamma-ray or cosmic-ray detector arrays on the ground (here mainly arrays of imaging atmospheric Cherenkov telescopes), the atmosphere enters in several ways: in the development of the particle showers, in the emission of light by shower particles, and in the propagation of Cherenkov light (or fluorescence light or of particles) down to ground level. Relevant parameters and their typical impact on energy scale and so on are discussed here. } ",
        "introduction": "Ground-based gamma-ray and cosmic-ray detectors make use of the atmosphere as a part of their instruments. Incoming energetic particles (neglecting neutrinos or exotic particles for the purpose of this paper) will interact with the atmosphere, produce secondary particles and initiate a particle cascade or extensive air shower. The level of inclusion of the atmosphere into the instrument differs between different detection techniques. Particle detector arrays are affected by the development of extensive air showers. Beyond the atmospheric overburden on ground, the atmospheric profile will also matter, due to the competition between interactions and decays of secondary unstable particles produced in interactions. Since most of the development of air showers happens at altitudes of up to a few ten kilometers of altitude, the change of atmospheric composition beyond about 80 km altitude has, fortunately, no significant impact onto the air shower development and detection.  For atmospheric Cherenkov detectors as well as for air fluorescence detectors, the atmosphere enters at two additional levels. The first is the emission process and the second is the propagation of the emitted light down to the photo-detectors. The Cherenkov and fluorescence processes depend in different ways on changes in density profiles. For the fluorescence technique, there is a competition between fluorescence light emission and collisional de-excitation, after the nitrogen molecules get excited by the charged particles in the air shower. The higher efficiency of light emission is balanced basically by lower excitation in lower density air. As a consequence, the fluorescence process depends relatively weakly on the air density profile, except for the shower development itself. The Cherenkov emission process is more strongly impacted by the density profile. Higher density, resulting in an almost proportionally enhanced refractivity (index of refraction minus one, $n-1$), results in enhanced emission as well as lowering the minimum particle velocity necessary for emission. It also increases the Cherenkov emission angle and spreading out the light on the ground over a larger area.  Both the Cherenkov and fluorescence techniques are affected by the light propagation, by absorption as well as by scattering. The absorption processes are not just affected by the main composition components but also affected by trace gases, like ozone or water vapour. At that stage the detailed composition profile thus enters the game. Scattering includes Rayleigh scattering - thus directly linked to the density profile - and Mie scattering on aerosols (also resulting in light absorption). As a consequence, these detection techniques are affected by the aerosol density, chemical composition, size distributions -- and all of these depend on altitude and change with time.  Simulations of air showers and detector response usually need some simplifications, in order to run in an efficient way. This paper will also try to point out some often-used simplifications which should be applied with care.  Since this paper is focusing on the impact of atmospheric parameters, a number of other site-related topics will be omitted, including the impact of the night-sky background (NSB) on Cherenkov or fluorescence detectors and the impact of the geomagnetic field. ",
        "conclusions": "Knowledge of the atmospheric density profile is important for proper simulation of any type of ground-based air shower instrument, for particle detector arrays as well as for Cherenkov or fluorescence detectors. For the latter two the profile enters also in the light emission processes. It is particularly relevant for the Cherenkov emission where the related index of refraction determines the minimum velocity for Cherenkov emission, the number of photons emitted per unit path length and also the Cherenkov cone opening angle. Most simulations assume a wavelength-independent index of refraction but the wavelength dependence can be turned on, at the expense of the simulation efficiency.  The extinction of Cherenkov and fluorescence light depends in addition to Rayleigh scattering and absorption on O$_2$ also on aerosols and on trace gases like ozone and water vapour. The tropospheric ozone absorption is of particular concern for Cherenkov detectors with extended UV response using muon rings for calibration. Hardly any simulations make use of measured ozone profiles. Except for this ozone part, the atmospheric transmission tables used in the simulations -- in particular the total aerosol optical depth -- can be best checked with star light extinction.  For further improving the accuracy of the simulations, the aerosol profile along the line of sight should be monitored. Nevertheless tricky are aerosol layers intersecting the shower development, requiring dedicated and well adapted simulations for correcting their effect.  Scattered light, in particular scattered Cherenkov light, is highly relevant to the fluorescence technique but only has a small effect for IACTs. For IACT simulations, it is generally ignored."
    },
    "1402/1402.0989_arXiv.txt": {
        "abstract": "Spectral and Temporal properties of black hole candidates can be explained reasonably well using Chakrabarti-Titarchuk solution of two component advective flow (TCAF). This model requires two accretion rates, namely, the Keplerian disk accretion rate and the halo accretion rate, the latter being composed of a sub-Keplerian, low angular momentum flow which may or may not develop a shock. In this solution, the relevant parameter is the relative importance of the halo (which creates the Compton cloud region) rate with respect to the Keplerian disk rate (soft photon source). Though this model has been used earlier to manually fit data of several black hole candidates quite satisfactorily, for the first time, we made it user friendly by implementing it into XSPEC software of GSFC/NASA. This enables any user to extract physical parameters of the accretion flows, such as two accretion rates, the shock location, the shock strength etc. for any black hole candidate. We provide some examples of fitting a few cases using this model. Most importantly, unlike any other model, we show that TCAF is capable of predicting timing properties from the spectral fits, since in TCAF, a shock is responsible for deciding spectral slopes as well as QPO frequencies. ",
        "introduction": "Compact objects, such as black holes, neutron stars, etc. are identified by electromagnetic radiations emitted from the accreting matter. Understanding the spectral and timing properties of this radiation is essential for model builders and theorists alike. In a binary system, matter from the companion star accrets into the black hole through the Roche lobe, and/or through capturing its winds. This matter produces a disk like structure around the primary compact object. There are a large number of theoretical models in the literature which explain the physics of accretion around a black hole. Evidences of the standard disk proposed by \\citet{SS73} and \\citet{NT73} % are present in most of the binary systems. However, the emitted spectrum of the radiation is multi-color in nature and contains both thermal and non-thermal components and a standard disk cannot explain the entire X-ray spectral features. Moreover, the inner region of the standard disk may be unstable due to the viscous and thermal effects \\citep{Lightman74,Kobayashi03}. %  Simply put, one of the components of the spectrum is a multi-color blackbody radiation from the standard Keplerian disk and the other is a power-law component formed due to repeated Compton scatterings of the low energy (soft) photons of this blackbody by the hot electrons of the `Compton' cloud \\citep{ST80,ST85}. % There are many speculations regarding the nature of this Compton cloud ranging from a magnetic corona \\citep{Galeev79}, to hot gas corona over the disk \\citep{Haardt93,Zdziarski03}. Since the formation process of a static corona around an accretion disk is totally unknown, and since a low angular dynamic flow may naturally act as a corona, \\citet[][hereafter CT95]{CT95} % proposed that a disk having two distinct components, a Keplerian disk submerged inside a sub-Keplerian halo is enough to explain all the spectral properties very satisfactory. Observational evidences also started to support this so-called two-component advective flow (TCAF) \\citep[e.g.,][]{Soria01,Smith02,Wu02,Cambier13}. % While creating a self-consistent TCAF solution, the properties of a viscous transonic flow was made use of, in which a flow having viscosity above a critical value naturally forms a Keplerian disk and the region with a lower viscosity, due to centrifugal barrier, forms a shock wave, typically, at a few tens of Schwarzschild radii. The post-shock region (from the shock and the inner sonic point) basically evaporates the Keplerian component and together acts as a Compton cloud which produces a power-law component (hard photons) with exponential cut-off in the spectrum through thermal Comptonization. From the inner sonic point to the horizon of the black hole (bulk motion dominated advective flow or BDAF) the matter is advected rapidly to the black hole. The bulk motion in this region also up-scatters the soft photons and produces a second power-law component even when the temperature of the region is zero. If the centrifugal barrier is not strong enough, the shock may not form, but the flow still slows down. The spectral properties in this case are discussed in \\citet[][hereafter C97]{C97}. The CENtrifugal pressure supported BOundary Layer or CENBOL, referred to the post-shock region or centrifugal force dominated region, which is also the base of the outflows where the pre-Jet is launched, plays the most important role in the black hole physics. As usual, this CENBOL, pre-Jet and BDAF intercept soft photons from the Keplerian disk and reprocess them to high energies via inverse Compton scattering. In this {\\it letter}, we will implement the TCAF solution to study spectral properties of black hole candidates (BHCs) using widely used user-friendly spectral analysis software package, developed by GSFC/NASA, called XSPEC. For the sake of concreteness, we focus only the cases where only the CENBOL is present. We ignore the effects of BDAF and pre-Jet. In our next version of analysis, these components and spin of the black hole would be included.  The Galactic transient black hole candidates are very interesting objects to study in X-rays because these sources generally show rapid evolutions in their temporal and spectral properties during their outburst phases, which are strongly correlated to each other \\citep[see for a review,][]{RM06}. In general, four basic states - {\\it hard, hard-intermediate, soft-intermediate}, and {\\it soft} states are observed during an outburst of the BHCs \\citep[see,][and references therein]{Nandi12}. The evolutions of these spectral states are observed, which indeed form a hysteresis loop during the outburst with hard states are found to be at the beginning and end time of the outbursts, whereas soft and intermediate spectral states are observed in between. The evolution of spectral states are strongly dependent on the variation of the accretion rates. According to the TCAF solution, accretion flow rates may be controlled by a physical parameter, such as the magnetic viscosity, perhaps owing to the enhanced magnetic activity of the companion \\citep{Wu02,Nandi12,DD13}. During the rising phase of the outburst, viscosity may cause an increase in the accretion rate of the Keplerian matter. As the viscosity is reduced, the Keplerian rate is reduced and declining phase starts. The Keplerian disk itself recedes away leaving behind only the low-angular sub-Keplerian flow causing a hard state. Thus, a rigorous fit with TCAF model is expected to throw light on how the accretion rates and the flow geometry evolve with time.  In general, low and intermediate frequency quasi-periodic oscillations (LFQPOs) are observed in hard and intermediate (hard-intermediate and soft-intermediate) spectral states of transient black hole candidates. These QPOs are reported extensively in literature, although still there are debates on the origin of these QPOs. However, according to the shock oscillation model (SOM) by Chakrabarti and his collaborators, LFQPOs are originated due to the oscillation of the post-shock region (\\citealt{MSC96}, hereafter MSC96; \\citealt{CAM04}, hereafter CAM04; \\citealt{GGC14}, hereafter GGC14) when the resonance occurs between the infall time scale and the cooling time scale in CENBOL. During oscillation, the shape of the Compton cloud and the degree of interception change periodically. Since from our spectral fit, we can directly extract values of physical parameters related to this shock wave, we can also predict what should be the frequency of the observed low frequency QPO (if present; see \\S 4.1 for details).  The {\\it paper} is organized in the following way: in the next Section, we briefly describe properties of the TCAF model. In \\S 3, we discuss the method of the implementation of the TCAF model in XSPEC for spectral fittings. In \\S 4, TCAF model fitted results obtained from the spectral fit of three different BHCs. Finally, in \\S 5, we make concluding remarks and our future work plans. ",
        "conclusions": ""
    },
    "1402/1402.1177_arXiv.txt": {
        "abstract": "{ Polaris the Cepheid has been observed for centuries, presenting surprises and changing our view of Cepheids and stellar astrophysics, in general. Specifically, understanding Polaris helps anchor the Cepheid Leavitt law, but the distance must be measured precisely.  The recent debate regarding the distance to Polaris has raised questions  about its role in calibrating the Leavitt law and even its evolutionary status.  In this work, I present new stellar evolution models of Cepheids to compare with previously measured CNO abundances, period change and angular diameter.  Based on the comparison, I show that Polaris cannot be evolving along the first crossing of the Cepheid instability strip and cannot have evolved from a rapidly-rotating main sequence star.  As such, Polaris must also be at least 118~pc away and pulsates in the first overtone, disagreeing with the recent results of Turner et al. (2013).  } ",
        "introduction": "The North Star, Polaris (HD~8890), has fascinated people for centuries, guided explorers, and appears in the mythologies of numerous cultures.  Polaris is also the nearest classical Cepheid, making it a powerful laboratory for understanding stellar evolution as well as anchoring the calibration of the Cepheid Leavitt Law (period-luminosity relation).  However, one of the greatest hindrances to understanding Polaris is its distance.  \\cite{vanLeeuwen2007} measured a distance $d = 129 \\pm 2~$pc from revised Hipparcos parallaxes, whereas \\cite{Turner2013} found a distance $d = 99\\pm 2$~pc based on spectroscopic line ratio measurements.  The two distances are both measured precisely, yet disagree significantly.  Furthermore, these distances are particularly interesting because they suggest different pulsation properties and evolutionary histories for Polaris.  Analysis of interferometric observations suggest an angular diameter for Polaris of $\\theta = 3.123 \\pm 0.008$~mas \\citep{Merand2007}, hence a mean radius of either $33.4 \\pm 0.6~R_\\odot$ or $43.5 \\pm 0.8~R_\\odot$ depending on which distance one considers.  If the radius is the former value then the period-radius relation \\citep[e.g.][]{Gieren1997, Neilson2010, Storm2011} implies that Polaris is pulsating as a fundamental mode Cepheid, whereas if the radius is the larger value then Polaris must be pulsating as a first overtone Cepheid.  Understanding the pulsation mode is necessary for calibrating the Cepheid Leavitt Law \\citep{vanLeeuwen2007} and constraining the structure of the Cepheid instability strip on the Hertzsprung-Russell diagram.  The distance debate also presents challenges for understanding the evolution of Polaris.  Assuming Polaris is at the closer distance, \\cite{Turner2013} argued that Polaris is evolving along the first crossing of the Cepheid instability strip because of its large rate of period change \\citep{Turner2006, Neilson2012a}.  In that case, Polaris has evolved from a main sequence star but has not yet become a red giant star.   \\cite{Turner2013} also suggested that Polaris must have rotated rapidly as a main sequence star to explain the measured nitrogen and carbon abundances \\citep{Usenko2005}.  On the other hand, \\cite{Neilson2012a} suggested that Polaris cannot be evolving along the first crossing if it is at the farther distance because the measured rate of period change is too small relative to a first-crossing Cepheid with a similar luminosity. Hence, Polaris is evolving along the third crossing of the instability strip on the Cepheid blue loop and its observed CNO abundances are consistent with dredge up during the previous red giant stage. \\cite{Neilson2012a} further suggested that to account for discrepancy between measured and predicted rates of period change, Polaris must be losing mass in an enhanced stellar wind with $\\dot{M} = 10^{-7}$ -- $10^{-6}~M_\\odot$~yr$^{-1}$.  Our understanding of Polaris is biased by the assumed distance, therefore measuring a precise and consistent distance is crucial.  However, there is no a priori reason to prefer one measured distance over another.  \\cite{Turner2013} argued that Hipparcos parallax must be wrong because it  is inconsistent with other measurements such as space velocities of nearby stars, and main sequence fitting of Polaris's F3~V binary companion.  \\cite{vanLeeuwen2013} countered that the first argument could not be reproduced, suggesting that Polaris is not a member of the local cluster. Furthermore, the F3~V stars have a large enough color distribution that would be consistent with both distances.  \\begin{figure*}[th] \\begin{center} \\includegraphics[width=0.48\\textwidth]{plot1-9.eps}\\includegraphics[width=0.48\\textwidth]{plot1-6.eps} \\end{center} \\caption{Predicted rates of period change for stars crossing the Cepheid instability strip assuming different prescriptions for Cepheid mass loss: $\\dot{M} = 10^{-9}~M_\\odot~$yr$^{-1}$ (left) and $10^{-6}~M_\\odot~$yr$^{-1}$ (right), the results for the other two cases are nearly indistinguishable.  The grey band refers to the measured rate of period change for Polaris \\citep{Neilson2012a}.} \\label{plot_pdot} \\end{figure*}  There are also questions regarding the spectroscopic method employed by \\cite{Turner2013}. That method is based on the FGK supergiant calibration of \\cite{Kovtyukh2007} and \\cite{Kovtyukh2010}, hence, stellar effective temperatures and luminosities are determined by measuring various iron line ratios from optical spectra.  However, Cepheid variable stars are not static yellow supergiants. For instance, \\cite{Kervella2006} and \\cite{Merand2006, Merand2007} detected infrared excesses about a sample of Galactic Cepheids, but not for the yellow supergiant star $\\alpha$ Persei. There are also observations of dynamic motions in the photosphere driven by pulsation that do not occur in static yellow supergiant stars \\citep[e.g.][]{Nardetto2008}. \\cite{Engle2012} reported X-ray and UV observations of the nearest Cepheids including Polaris, suggesting the presence of hot ($\\approx 100,000$~K) plasma in the photosphere, which may affect the ionization rates of various species.   These differences could affect the systematic uncertainties of the effective temperature and luminosity measured for Polaris.  However, in spite of these differences, \\cite{Turner2013} reported a mean effective temperature $\\langle T_{\\rm{eff}} \\rangle = 6025\\pm 1$~K and $\\langle M_V \\rangle = 3.07\\pm 0.01$, where the presented error is statistical only.  There are no systemic uncertainties presented.  Convection alone would lead to a greater variation of temperature. However, to be consistent with the Hipparcos parallax the uncertainty of the luminosity must be $\\ge 0.6$ dex.  The current distance debate to Polaris is confusing our understanding of the star from the context of stellar evolution and pulsation.  In this article, I compute new state-of-the-art stellar evolution models to compare with various observations of Polaris, such as the rate of period change \\citep{Turner2005, Neilson2012a}, angular diameter \\citep{Merand2006}, and CNO abundances \\citep{Usenko2005}.  We will also compare these models to the estimates of the distance and effective temperature, hence radius and luminosity.  Comparing stellar evolution models to the $[N/H]$ and $[C/H]$ abundances test the hypothesis that Polaris is evolving along the first crossing of the instability strip, i.e., Polaris was once a rapidly rotating during main sequence star.   Comparing models with the observed rate of period change provides insight into the mass-loss rate and which crossing of the instability strip Polaris is evolving, while comparing the models to the angular diameter constrains the radius and luminosity of the star, hence its distance.  In Sect.~2, I briefly describe the stellar evolution models computed. In Sect.~3, I describe how Cepheid rates of period change are computed from stellar evolution models. I also compute rates of period change as a function of stellar mass-loss rate to compare to the observed rate of period change.  I compare stellar evolution models with the observed nitrogen and carbon abundance to constrain the hypothesis that Polaris is a first-crossing Cepheid in Sect.~4.  In Sect.~5, I discuss the implications of these tests for the distance and pulsation mode. ",
        "conclusions": "The debate over the distance to Polaris is important for the calibration of the Cepheid Leavitt Law, as the nearest Cepheids will calibrate the zero-point \\citep[e.g.][]{Freedman2012}. Resolving the debate is also important for constraining stellar evolution physics, and constraining the properties of other Cepheids, which will become particularly important as GAIA observations will measure parallaxes for thousands of Cepheids \\citep{Windmark2011}.  In this work, I presented new stellar evolution models to reanalyze the fundamental properties of Polaris, this time without assuming a distance.  The observed rate of period change is found to be inconsistent with stellar evolution models for Cepheids, for both the first and third crossings of the instability strip, regardless of whether the \\cite{vanLeeuwen2013} or \\cite{Turner2013} distance is assumed.  There is no obvious theory where Polaris is consistent with stellar evolution on the first crossing of the instability strip, since that time scale is determined by dynamical time scales.   On the other hand, mass loss could explain the rate of period change for a Cepheid on the third crossing, even though mass loss also reduces the effective temperature width of the Cepheid blue loop making it difficult to compare with the observed properties of Polaris.  However, it can be noted that the measured period change is inconsistent with a first-crossing Cepheid.  I further computed new rotating stellar evolution models to explore the hypothesis that Polaris is a first-crossing Cepheid with nitrogen and carbon abundances changed by rotational mixing.  Rotational mixing models predict abundances roughly consistent with measurements by \\cite{Usenko2005} during the first-crossing of the instability strip, however, the predicted rotation rates are too large to be consistent with the measured turbulent velocity for Polaris \\citep{Usenko2005} and rotational velocities for a sample of Galactic Cepheids \\citep{Bersier1996}. As such, the properties of Polaris are inconsistent with being a first-crossing Cepheid and must be a third-crossing Cepheid. Because Polaris is a third-crossing Cepheid, then it must also be at a distance greater than $118$~pc, significantly greater than that measured by \\cite{Turner2013} and must also be pulsating as a first-overtone Cepheid.  While the results of \\cite{Turner2013} are inconsistent with stellar evolution calculations, there is still much to explore about the nearest Cepheid, Polaris to constrain the physics behind stellar multiplicity \\citep[e.g.][]{Evans2013}, mass loss \\citep{Neilson2012a, Neilson2012b}, ultraviolet and x-ray radiation \\citep{Engle2012}, the apparent brightening of Polaris \\citep{Engle2014} and late-stage evolution \\citep[e.g.][]{Langer2012} as we approach an era of precision stellar astrophysics."
    },
    "1402/1402.4108_arXiv.txt": {
        "abstract": "In this follow-up work to the High Energy Physics Community Summer Study 2013 (HEP CSS 2013, a.k.a. \\textsc{Snowmass}), we explore the scientific capabilities of a future Stage-IV Cosmic Microwave Background polarization experiment (CMB-S4) under various assumptions on detector count, resolution, and sky coverage.  We use the Fisher matrix technique to calculate the expected uncertainties in cosmological parameters in $\\nu\\Lambda$CDM that are especially relevant to the physics of fundamental interactions, including neutrino masses, effective number of relativistic species, dark-energy equation of state, dark-matter annihilation, and inflationary parameters.  To further chart the landscape of future cosmology probes, we include forecasted results from the Baryon Acoustic Oscillation (BAO) signal as measured by DESI to constrain parameters that would benefit from low redshift information. We find the following best 1-$\\sigma$ constraints: $\\sigma(M_{\\nu})$ $= 15$ meV, $\\sigma(N_{\\rm eff})$ $= 0.0156$, Dark energy Figure of Merit = 303, $\\sigma(p_{ann})$ $= 0.00588\\times3\\times10^{-26}$ cm$^3$/s/GeV,  $\\sigma(\\Omega_K)$ $= 0.00074$,  $\\sigma(n_s)$ $= 0.00110$,  $\\sigma(\\alpha_s)$ $=  0.00145$, and  $\\sigma(r)$ $= 0.00009$. We also detail the dependences of the parameter constraints on detector count, resolution, and sky coverage. ",
        "introduction": "In the past two decades, CMB experiments have made great strides in sensitivity improvement. For satellite experiments, there was a factor of 10 better in sensitivity from \\textit{COBE} to \\textit{WMAP} and from \\textit{WMAP} to \\textit{Planck}~\\cite{2013arXiv1309.5383A}. These satellites, along with numerous ground-based and balloon-borne experiments, as listed in Ref. \\cite{Lambda}, collectively moved us from setting upper limits on the temperature anisotropy ($C_\\ell^{TT}$), the E- and B-mode polarization ($C_\\ell^{EE}$ and $C_\\ell^{BB}$) spectra to precision measurements of the temperature power spectrum (e.g.~\\cite{2013ApJ...779...86S, 2013arXiv1303.5075P, 2013arXiv1301.1037D}), shrinking error bars on the E-mode polarization spectrum (e.g.~\\cite{2012ApJ...760..145Q, 2013arXiv1310.1422B}), and detecting gravitational lensing in the B-mode polarization power spectrum~\\cite{2013arXiv1307.5830H, 2013arXiv1312.6645P,2013arXiv1312.6646P}.  Tremendous opportunities still lie in CMB polarization. CMB polarization maps can be projected into a curl-free component (E-modes), and a divergence-free component (B-modes). Scalar and tensor perturbations seeded the temperature anisotropies in the CMB. While E-modes, like temperature anisotropies, could be sourced by both scalar and tensor perturbations through Thomson scattering, the only primordial source for B-modes is from tensor perturbations i.e. gravitational waves. Detecting this B-mode polarization signal, which peaks at the degree angular scale, would be a powerful probe of inflation.  B-modes at smaller angular scales are generated by gravitational lensing of E-modes by Large-Scale Structures (LSS). The conversion from E- to B-modes provides an exceptionally clean way of reconstructing the lensing potential $\\phi$ and its power spectrum $C_\\ell^{\\phi\\phi}$ \\cite{2003PhRvD..67h3002O}. The latter is sensitive to the properties of the LSS, making it a probe of the late universe ($z < 10$ \\cite{2006PhR...429....1L}). As a result, among other astrophysical properties, it is central to the measurements of the total neutrino mass and dark-energy equation of state.  Different techniques have been developed to reconstruct lensing potential from CMB temperature and polarization maps \\cite{2003PhRvD..67h3002O, Hirata:2003ka}. Recently, experiments have detected gravitational lensing in CMB temperature maps with high enough significance to measure the lensing potential power spectrum ~\\cite{2011PhRvL.107b1301D,2012ApJ...756..142V,2013arXiv1303.5077P}. Moreover, the first reconstruction through E-to-B channel has recently been demonstrated \\cite{2013arXiv1312.6646P}.  The lensing B-mode signal also acts as a foreground to the inflationary B-mode signal. Specifically, if the tensor-to-scalar ratio $r$, which parametrizes the amplitude of the tensor perturbation, is smaller than 0.02, the level of inflationary B-modes would be lower than the lensing B-modes \\cite{Kesden:2002ku,Knox:2002pe}. Algorithms have been developed to remove the lensing-induced B-modes \\cite{2012JCAP...06..014S} to study the primordial signal. Therefore, measuring the lensing potential to high precision is fundamental to advancing our understanding of both the early and the late universe.  We envision the Stage-IV\\footnotemark  \\footnotetext{Each advance in stage refers roughly to an order of magnitude increase in number of detectors \\cite{2013arXiv1309.5383A}.} Cosmic Microwave Background polarization experiment (CMB-S4) to be a high-resolution, high-sensitivity experiment for this very purpose. This paper explores how much CMB-S4 could help answer the pressing questions in High Energy Physics (HEP)  --- more broadly, physics of fundamental interactions. In what follows, we outline the main HEP topics that can be addressed by CMB polarization.\\\\  \\textbf{Cosmic Neutrino background --} The relic neutrino number density for each species is 112 cm$^{-3}$ \\cite{Fukugita}. This makes neutrinos the second most abundant particle in the universe, after photons. Although neutrinos are massless in the Standard Model, solar and atmospheric neutrino oscillation experiments have shown that they are massive. Assuming normal hierarchy, a lower limit on the total neutrino mass is $\\sim58$ meV. Though minuscule in mass, their large cosmic abundance enables us to observe their cumulative gravitational effect on structure formation: neutrinos suppress structure growth below scales defined by their free streaming distances. As a result, we can measure the total neutrino mass by observing the matter power spectrum -- which CMB lensing is particularly sensitive to~\\cite{2003PhRvL..91x1301K}. Another way to study the cosmic neutrino background is to measure the extra number of relativistic species $N_{\\text{eff}}$ aside from photons at recombination. In the Standard Model, no other particles except neutrinos contribute as extra relativistic species. In this scenario, $N_{\\text{eff}}$ is predicted to be 3.046. Deviation from this value hints at new physics. The CMB is uniquely sensitive to the photon-baryon-dark matter interaction at recombination. The current constraints on $N_{\\text{eff}}$ are  $3.30^{+0.54}_{-0.51}$ (at the $95\\%$ C.L.) from \\textit{Planck} + \\textit{WMAP} polarization + small scale CMB data + BAO measurements, and the current limit on total neutrino mass is $< 0.23\\, \\text{eV}$ (at the $95\\%$ C.L.) from \\textit{Planck} data + BAO data \\cite{2013arXiv1303.5076P}. \\\\  \\textbf{Dark Energy --} Since the discovery of dark energy \\cite{1998AJ....116.1009R,1999ApJ...517..565P}, its contribution to the energy density of the universe $\\Omega_{\\Lambda}$ has been determined to percent level over the past 15 years.  However, its fundamental nature still remains mysterious. The three broad possibilities are: (1) vacuum energy manifesting as a cosmological constant; (2) a spatially homogeneous dynamical field; (3) modified gravity on cosmological scales.  These hypotheses can be tested experimentally by measuring the dark-energy equation of state $w$ through the clustering of LSS as a function of redshift. So far, measurements are consistent with a cosmological constant.  Over the next decade, photometric and spectroscopic surveys such as LSST, Euclid, and DESI will take these tests to the next level of precision, opening up opportunities for new discoveries \\cite{Albrecht2012}.  CMB lensing is highly complementary to these surveys, because it provides high-redshift information in the linear regime, cf.~\\cite{2001PhR...340..291B}.  It is especially valuable for searches of early evolution of $w$ when $z>2$, which would not be possible with galaxy surveys. In addition, CMB lensing can help optical lensing surveys calibrate their shear bias and boost constraints on cosmological parameters, \\cite{2013arXiv1311.0905P, 2013arXiv1311.2338D, 2013arXiv1311.6200H}. The current constraints of $w_0$ and $w_a$ in the $w = w_0+w_a(1-a)$ model are: $w_0 = -1.04^{+0.72}_{-0.69}$ and $w_a < 1.32$  (\\textit{Planck}+\\textit{WMAP} polarization + BAO data), cf.~\\citet{2013arXiv1303.5076P}. \\\\  \\textbf{Dark-matter annihilation--} The gravitational properties of dark matter (DM) have been well constrained by data from a plethora of measurements, e.g. weak~\\cite{2003ARA&A..41..645R} and strong~\\cite{1998ApJ...498L.107T} lensing, multi-wavelength studies of the Bullet Cluster \\cite{2006ApJ...648L.109C}, distant supernovae~\\cite{1998AJ....116.1009R,1999ApJ...517..565P}, and the CMB~\\cite{2013arXiv1303.5076P}. However, the particle nature of dark matter remains unknown. There are several types of experiments aiming to understand dark matter interactions: (1) direct detection experiments, which rely on nuclear recoil signature of DM-Standard Model particle scattering, (2) indirect detection experiments, which are sensitive to the products of DM decay or annihilation, and (3) collider experiments. For a DM theory and detection review, see Ref. \\cite{2010ARA&A..48..495F}.  An alternative observable is the CMB two-point correlation function. When DM annihilates, heat is transferred to the photon-baryon fluid, atoms are ionized and/or excited~\\cite{2004PhRvD..70d3502C}. \\citet{2013PhRvD..87l3513S} performed a detailed study regarding the energy deposition history into the photon-baryon fluid.  Dark-matter annihilation also leads to growing ionization fraction perturbations and amplified small-scale cosmological perturbations, leaving an imprint on the CMB bispectrum \\cite{2013PhRvD..87j3522D}. CMB temperature and polarization spectra can constrain the parameter $p_{ann} \\equiv  f \\langle \\sigma v \\rangle\\,/\\,m_{DM}$, where $f$ is the fraction of energy deposited into the plasma, $\\langle \\sigma v \\rangle$ is velocity-weighted cross section, and $m_{DM}$ is the mass of the DM particle. Current constraints coming from \\textit{WMAP} $9$-year data, \\textit{Planck}, ACT, SPT,  BAO, HST and SN data exclude DM masses below $26$ GeV at the 2$\\sigma$ level, assuming that all the energy is deposited in the plasma \\cite{2013arXiv1310.3815M}. We show in this paper that CMB-S4 will tighten these constraints by a factor of $10$.\\\\  \\textbf{Inflation --} A detection of degree-scale B-mode polarization would help us learn about the physics of inflation. In particular, the amplitude of tensor perturbations is directly related to the energy scale of inflation. In addition, whether $r$ is above or below $\\sim$0.01 determines whether the inflaton field range is sub- or super-Planckian in a broad class of inflationary models. Beyond B-modes, small angular scale E-mode polarization, which is less contaminated by foregrounds than the temperature maps, can test the following predictions from slow-roll inflation: (1) the nearly scalar invariant primordial power spectrum where $n_s$ is close to but not exactly 1, (2) the small running of $n_s$, and (3) the almost flat mean spatial curvature, $\\Omega_K  \\approx 10^{-4}$. The current tightest constraint for $n_s$ is $0.9603 \\pm 0.0073$ at the $68\\%$ C.L., for $dn_s/d\\log k$ is $-0.014^{+0.016}_{-0.017}$ at the $95\\%$ C.L., and for $\\Omega_K$ is $-0.0005^{+0.0065}_{-0.0066}$ at the $95\\%$ C.L. \\cite{2013arXiv1303.5076P}. \\\\  We show in this paper, given the range of possible experimental configurations (survey coverage, depth, resolution), how sensitive future ground-based experiments like CMB-S4 will be in probing these areas of new physics. We deliberately extend the experimental configuration space stated in the Snowmass study for CMB-S4 so that we understand the steepness of the improvement in parameter constraints as a function of experimental setups. Furthermore, we detail the forecast methods which are lacking in the Snowmass report.  The paper is organized as follows. First, in Section~\\ref{sec:methods}, we describe the methodology and the fiducial cosmology.  Sections ~\\ref{sec:neutrinos} to \\ref{sec:Inflation} are assigned to the aforementioned HEP topics and parameters of interests: \\begin{itemize} \\item Section~\\ref{sec:neutrinos}: total neutrino mass and extra relativistic species -- $M_{\\nu}$ and  $N_{\\rm{eff}}$ \\item Section~\\ref{sec:DarkEnergy}: dark-energy equation of state -- $w_0$ and $w_a$ \\item Section~\\ref{sec:DarkMatter}: dark-matter annihilation -- $p_{ann}$ \\item Section~\\ref{sec:Inflation}: inflation -- $\\Omega_K$, $n_s$, $\\alpha_s$ (running of $n_s$), $r$, and $n_t$ \\end{itemize}  For each section above, we present a brief overview of the underlying physical phenomena and their effects on the CMB power spectra,  and present and discuss the results of the forecast. For the tensor-to-scalar ratio $r$, we have a more thorough method section outlined in Section~\\ref{sssec:r-method}.  Finally, we conclude and summarize the main results in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this work, we forecasted how well a highly capable next-generation ground-based CMB experiment can constrain cosmological parameters of fundamental physics relevant to both high energy physics and cosmology. We forecast for a range of experimental inputs -- $10^4$ to $10^6$ detectors, $1'-4'$ beams ($6'$, $8'$ in some cases), $1-75\\%$ sky fraction in order to see how the constraints for each parameters vary with these inputs.  We detailed in section~\\ref{sec:methods} the methods used to estimate the performance of a given CMB experimental design. We presented our results in sections~\\ref{sec:neutrinos},~\\ref{sec:DarkEnergy}, ~\\ref{sec:DarkMatter}, and ~\\ref{sec:Inflation}. Here we quote the range of constraints each parameter falls in for experiments with $10^4-10^6$ detectors, $1'$ to $4'$ beam, and 25\\% to 75\\% $f_{sky}$ as illustrations and summarize the CMB only parameter improvement dependence on $N_{\\rm det}$, beam size, and $f_{sky}$: \\begin{itemize} \\item $ 0.0156  \\leq \\sigma(N_{\\rm eff}) \\leq 0.0690$ (CMB): \\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $\\sigma(N_{\\rm eff})$ by about 30\\%, same for going from $10^5$ to $10^6$; decreasing beam size improves $\\sigma(N_{\\rm eff})$ by $\\sim$10\\% per arc-minute; $\\sigma(N_{\\rm eff})$ is not sample variance limited.  \\item $ 15  \\leq \\sigma(M_{\\nu}) \\leq 24 $ [meV] (CMB+BAO): \\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $ \\sigma(M_{\\nu})$ by 10-20\\% (smaller beam gives better improvement with $N_{\\rm det}$), while increasing $N_{\\rm det}$ from $10^5$ to $10^6$ improves $\\sigma(M_{\\nu})$ by 5-15\\%; decreasing beam size improves $\\sigma(M_{\\nu})$ at percent-levels per arc-minute; $\\sigma(M_{\\nu})$ is sample variance limited.  \\item $164 \\leq$ DETF-FoM $\\leq 303$ (CMB+BAO+$H_0$):\\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves FoM by more than factor of 2, same for going from $10^5$ to $10^6$; decreasing beam size improves the FoM by a few \\% to 10s of \\% depending on configuration; FoM is not sample variance limited.  \\item $0.00588 \\leq \\sigma(p_{ann}) \\leq 0.0110\\, [3\\times10^{-26}$ cm$^3$/s/GeV] (CMB):\\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $\\sigma(p_{ann})$ by about 4\\%, same for going from $10^5$ to $10^6$; decreasing beam size improves $\\sigma(p_{ann})$ by $\\lesssim$ 1\\%; $\\sigma(p_{ann})$ is sample variance limited.  \\item $0.00074 \\leq \\sigma(\\Omega_K) \\leq 0.0014 $ (CMB+BAO+$H_0$):\\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $\\sigma(\\Omega_K)$ by about 20\\%, while increasing $N_{\\rm det}$ from $10^5$ to $10^6$ improves $\\sigma(\\Omega_K)$ by 10-20\\% (smaller beam gives a better improvement with $N_{\\rm det}$); decreasing beam size improves $\\sigma(\\Omega_K)$ at percent levels per arc-minute; $\\sigma(\\Omega_K)$ is sample variance limited.  \\item $0.00110 \\leq \\sigma(n_s) \\leq 0.00236$ (CMB+BAO): \\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $\\sigma(n_s)$ by $\\sim\\, 5-10\\%$ (smaller beam gives a better improvement with $N_{\\rm det}$), same for going from $10^5$ to $10^6$; decreasing beam size improves $\\sigma(n_s)$ at percent levels per arc-minute; $\\sigma(n_s)$ is sample variance limited.  \\item $0.00145 \\leq \\sigma(\\alpha_s) \\leq 0.00330$ (CMB+BAO): \\\\ increasing $N_{\\rm det}$ from $10^4$ to $10^5$ improves $\\sigma(\\alpha_s)$ by 16-20\\%, and by 13-17\\% going from $10^5$ to $10^6$ (smaller beam gives a better improvement with $N_{\\rm det}$); decreasing beam size improves $\\sigma(\\alpha_s)$ at percent levels per arc-minute; $\\sigma(\\alpha_s)$ is sample variance limited.  \\item $0.00009 \\leq \\sigma(r) \\leq 0.00203$ for 1\\% and 10\\% foreground residual: \\\\ $\\sigma(r)$ is foreground limited; when foreground is high, the optimal $f_{sky}$ shifts higher; decreasing beam size improves $\\sigma(r)$ at percent levels per arc-minute (the slope of this trend with beam increases with lower sky coverage).  \\end{itemize} Detailed constraints in specific cases can be read off from tables listed in each section.\\\\  Besides learning the approximate ranges of how well these parameters can be constrained with CMB-S4, we also learn how the constraints improve as functions of $N_{\\rm det}$, beam size, and $f_{sky}$: \\begin{itemize} \\item For all parameters, except $r$, increasing $f_{sky}$ always improve the constraints even though the overall sensitivity of the experiment decreases. \\item For all parameters, except those related to dark-energy equation of state, going from $10^5$ to $10^6$ detectors yields the same or less percentage improvement on the constraints than going from $10^4$ to $10^5$ detectors. The improvements range from a few to tens of percent. For $M_{\\nu}$, the improvement beyond $10^5$ detectors is marginal when BAO signal is added. \\item Dependence on beam size is quite mild -- constraints on $N_{\\rm eff}$ improves by about 10\\% per arc-minute decrease while for all other parameters it is around a few \\% improvement per arc-minute decrease. \\end{itemize}  We envision CMB-S4 to be a powerful next-generation ground-based CMB polarization experiment with high-resolution and high-sensitivity. With CMB-S4, we showed that most constraints on cosmological parameters are sample variance limited. Combining these data sets with space-borne observations will allow access to larger sky fraction, thus further improving the constraints on sample variance limited parameters."
    },
    "1402/1402.1207_arXiv.txt": {
        "abstract": "A high performance Space-Time Reference in orbit could be realized using a stable atomic clock in a precisely defined orbit and linking that to high accuracy atomic clocks on the ground using a laser based time-transfer link.  This would enhance performance of existing systems and provide unique capabilities in navigation, precise timing, earth sciences, geodesy and the same approach could provide a platform for testing fundamental physics in space. Precise laser time- and frequency-transfer from the ground to an orbiting satellite would make it possible to improve upon the current state of the art in timing (about 1 to 30\\,ns achieved with GPS) by roughly a factor of 1000 to the 1\\,ps level. ",
        "introduction": "There have been tremendous advances in the performance of atomic frequency standards (clocks) over the past 40 years, and, for compelling reasons, there are growing efforts to put more advanced atomic clocks into space.   Prominent examples are the PHARAO cold-cesium atomic clock that is part of the European ACES mission\\cite{ESA_site} scheduled to fly on the International Space Station about 2016, the compact Hg$^+$ ion standard of JPL designed for space applications, and other promising systems under development for the future (the ESA Space Optical Clock,\\cite{Schiller_SOC_project} the DARPA Slow Beam Optical Clock\\cite{DARPA_AOSense}). Advanced laser systems have already vastly improved the performance of atomic clocks and optical frequency synthesis and division. Lasers can do the same for time transfer to space. ",
        "conclusions": ""
    },
    "1402/1402.0751_arXiv.txt": {
        "abstract": "In order to build more realistic single stellar population (SSP) models with variable $\\alpha$-enhancement, we have recently determined [Mg/Fe] in a uniform scale with a precision of about 0.1 dex for 752 stars in the MILES empirical library. The [$\\alpha$/Fe] abundance ratio is commonly used as a good temporal scale indicator of star formation, taking Mg as a template for $\\alpha$ elements. Calcium is another element whose abundance is currently being investigated for the MILES stars. The MILES library is also being expanded by around 20\\% by including stars with known {\\it T}$_{\\rm eff}$, log~$g$, [Fe/H] and [Mg/Fe]. The transformation of their photospheric parameters to the MILES system has been carried out, but the calibration of their [Mg/Fe] is still in progress. In parallel, C, N and O abundances are also being compiled from literature for the library stars because they play an important role in the photospheric opacity, particularly influencing the blue spectral region. The Galactic kinematic classification of MILES stars with compiled [Mg/Fe] has been just computed such that this information can be considered in the SSP modelling. Comparisons of theoretical stellar predictions of the Lick line-strength indices against the MILES data have revealed the good behaviour of Fe-sensitive indices predictions, while highlighting areas for improvement in some models for the higher order H-Balmer features. ",
        "introduction": "\\label{s:intro}  The main limitation of current empirical stellar spectral libraries is the use of stars whose individual elemental abundances are not adequately considered. Moreover, the stars in these libraries are basically selected from the solar neighbourhood, in which the chemical evolution of Milk Way is imprinted. That limitation is also propagated to the SSP isochrone-based models that have been built on these libraries. Typically, only iron is taken as a metallicity tracer, but the spectral energy distributions of stars and stellar systems considerably depend on abundance ratios of other metals relative to iron (e.g. CNO group and alpha elements). For instance, [$\\alpha$/Fe] is known as a good temporal scale indicator for the star formation since the alpha and iron-group elements have distinct nucleosynthetic origins with different time scales, respectively type II and type Ia supernovae. The higher this ratio is, the shorter the time scale for stellar formation. Magnesium is usually assumed as a template for $\\alpha$ elements.  Therefore we have been taking steps to improve one of the most widely used empirical stellar spectral libraries - MILES \\citep{2006MNRAS.371..703S} \\& \\citep{2007MNRAS.374..664C}, and to use it to test theoretical models. Our ultimate aim is to utilise a combination of well characterised stars and theoretical model spectra to predict SSP spectra for a wide range of abundance patterns. Furthermore, building a new generation of semi-empirical SSP models with variable [$\\alpha$/Fe] will open new possibilities to minimise the effect of extrapolations of [$\\alpha$/Fe] in the models parametric space as well as to break the age-metallicity degeneracy over the model line-strength predictions, improving the recovery of star formation history in galaxies. ",
        "conclusions": "\\label{s:maths}  We have determined [Mg/Fe] with a precision of $\\sim$0.1 dex for about 80\\% of MILES stars that are placed on a uniform scale (available on request or in Milone et al. 2011). A robust spectroscopic analysis was carried out using the MILES spectra and LTE spectral synthesis of two Mg features. This same semi-automated approach can be applied to recover [Ca/Fe] and perhaps [C/Fe]. MILES is currently being expanded by around 20\\% through the inclusion of stars with known {\\it T}$_{\\rm eff}$, log~$g$, [Fe/H] and [Mg/Fe] that fills in some gaps of the library 4-D parameter space and increase the star density in other regions. The transformation of their photospheric parameters to the MILES homogeneous system has been carried out, and the calibration of their [Mg/Fe] will be completed soon. The abundances of calcium, carbon, nitrogen and oxygen are currently being investigated for the MILES stars (compilation from literature for all of them and Ca atomic features spectroscopically analysed at mid-resolution). Specifically carbon and nitrogen most likely play an important role in the blue spectral region as we have concluded from our empirical analysis based on MILES spectral ratios between pairs of similar stars in {\\it T}$_{\\rm eff}$ and log~$g$. The Galactic kinematic classification of MILES stars with compiled [Mg/Fe] has been just computed such that this information can be included in the SSP modelling. Careful comparisons of the theoretical stellar predictions of the Lick indices against the MILES data have revealed the good behaviour of Fe-sensitive indices predictions, while highlighting areas for improvement in some models for the higher order H-Balmer features. We intend to compute a set of self-consistent semi-empirical SSP models with variable $\\alpha$-enhancement for a range in ages and metallicities around solar."
    },
    "1402/1402.4458_arXiv.txt": {
        "abstract": "name}{Abstract} \\begin{abstract} In this work we focus on the evolution of the linear perturbations in the novel hybrid metric-Palatini theory achieved by adding a $f(\\mathcal{R})$ function to the gravitational action. Working in the Jordan frame, we derive the full set of linearized evolution equations for the perturbed potentials and present them in the Newtonian and synchronous gauges. We also derive the Poisson equation, and perform the evolution of the lensing potential, $\\Phi_{+}$, for a model with a background evolution indistinguishable from $\\Lambda$CDM. In order to do so, we introduce a designer approach that allows to retrieve a family of functions $f(\\mathcal{R})$ for which the effective equation of state is exactly $w_{\\textrm{eff}} = -1$. We conclude, for this particular model, that the main deviations from standard General Relativity and the Cosmological Constant model arise in the distant past, with an oscillatory signature in the ratio between the Newtonian potentials, $\\Phi$ and $\\Psi$. ",
        "introduction": "Introduction}  Einstein's General Relativity (GR) is modern cosmology's main framework, providing a set of equations that dictate the dynamics of our Universe according to its material constituents. By them, our Universe could be expanding, static, or even collapsing. However, it is now well established that our Universe is currently undergoing an accelerated expansion which was preceded by phases of matter and radiation domination where gravitational attraction resulted in a decelerated expansion. And, at the beginning, it should have experienced a period of quasi-exponential inflation, so that any primordial spatial curvature would have been wiped out, leading to the spatially-flat and homogeneous Universe we observe.  The simplest explanation for the Universe's accelerated expansion is a cosmological constant, $\\Lambda$, with a constant ratio of pressure to density (usually defined as the equation of state, $w$) equal to $-1$. Despite being in agreement with supernovae observations \\cite{accel1,accel2,accel3,accel4}, data from the cosmic microwave background (CMB) \\cite{cmb1,cmb2} including the recent {\\it Planck} data \\cite{planck1}, and large-scale structure (LSS) data \\cite{lss1}, cosmologists still struggle to account for the difference between the theoretically-expected value for its energy density and the observed one. If it exists, observationally, it should account for approximately $70 \\%$ of the Universe's total energy density, a value $121$ orders of magnitude smaller than that obtained from quantum field theory (for a review on $\\Lambda$, see Ref.~\\cite{lambdareview}).  In light of these issues, new physics may be in order to account for that major component of our Universe, usually labeled dark energy (DE). Some theories, such as quintessence, k-essence, and so on, propose scalar fields rolling in a potential (see Ref.~\\cite{quintessence} and references therein for a comprehensive review). Other theories consider higher dimensions, as in braneworld models such as the DGP model \\cite{DGP,branes}, or assume that GR fails on cosmological scales and propose corrections to Einstein's action. The latter are grouped as the so-called Modified Gravity Theories (MGT), such as the Brans--Dicke scalar--tensor theory \\cite{bd}, Galileon models \\cite{gall}, the Fab Four \\cite{fabfour}, $f(R)$ theories \\cite{frreview}, and many others. For an extensive review on MGT, see Ref.~\\cite{reviewall}.  In this paper, we focus on  a novel model, the hybrid metric-Palatini gravity \\cite{main1,main2}. In this type of theories, the usual Einstein-Hilbert Lagrangian is supplemented with an $f(\\mathcal{R})$ Palatini term. This type of hybrid theory arises when perturbative quantization methods are considered on Palatini backgrounds \\cite{quantization}, which are connected with non-perturbative quantum geometries \\cite{quantumgeo}.  Like the pure metric and Palatini cases, the hybrid theory has a dynamically equivalent scalar-tensor representation \\cite{main1,main2}. Those authors have also shown that the scalar field need not be massive in order to pass the stringent Solar System constraints \\cite{main1}, in contrast to the metric $f(R)$ theories, while possibly modifying the cosmological \\cite{hybridcosmo} and Galactic \\cite{hybridgala} dynamics due to its light, long-range interacting nature. In Ref. \\cite{hybridcosmo}, the criteria for obtaining cosmic acceleration was discussed. Alongside that, several cosmological solutions were derived, depending on the form of the effective scalar field potential, describing both accelerating and decelerating Universes.  In this work, we focus on cosmological perturbations in the hybrid metric-Palatini formalism in the Jordan frame. Therefore, in Section \\ref{I} we briefly introduce the hybrid metric-Palatini model and, in Section \\ref{designer}, we present the designer approach which allows us to retrieve a family of solutions for $f(\\mathcal{R})$ whose effective equation of state is $w_{\\textrm{eff}} = -1$. In Section \\ref{II} we derive the full set of perturbed cosmological equations and present them in the Newtonian and Synchronous gauges. Then, in Section \\ref{III} we derive the Poisson equation and re-express the perturbed potentials in terms of the lensing potential, $\\Phi_{+}$, and the slip, $\\chi$, which we numerically evolve. We finish in Section \\ref{conclusion} with the conclusions of this work. ",
        "conclusions": "Conclusion}  In this paper, we have derived the full set of perturbed Einstein equations for the novel hybrid metric-Palatini theory of gravity and presented them in the Newtonian and synchronous gauges. The latter, in particular, open the possibility of implementing this model in CAMB \\cite{camb} and give an in-depth analysis of the effects it can have at early times and even constrain its parameter space, which we leave for future work.  We have introduced a designer approach to obtain a family of functions $f(\\mathcal{R})$ that reproduce a cosmology indistinguishable from $\\Lambda$CDM, with an effective equation of state exactly equal to $w_{\\textrm{eff}} = -1$. This particular approach leads to models where the modifications from standard General Relativity are more significant in the distant past. And, even though one can tweak the free parameters to control such modifications in order that they are quite negligible at a redshift of $z_{i} \\approx 1000$, this can prove problematic at even earlier times if the departure from GR gets increasingly larger, as it seems to do. Potentially, one could observe an inversion in the sign of $G_{\\textrm{eff}}$, leading to an inversion of the effect of gravity. This was, however, avoided in our analysis. We would like to point out that, other background solutions were neglected, for now, due to the inability to consistently set the initial conditions for $F$ when $w_{\\textrm{eff}} \\neq -1$. This is being explored for future work.  We also derived the Poisson equation, which is substantially different from standard GR, given the inclusion of several extra dynamical elements. We then introduced the lensing potential, $\\Phi_{+}$, and the slip between the Newtonian potentials, $\\chi$, which we numerically evolved using the designer approach. We note that the departure from GR is more noticeable at the beginning of evolution. More specifically, $\\chi$ oscillates with a frequency proportional to the mode's wave number, and the oscillations' amplitude is the largest at this point. It is then gradually damped due to the evolution, tending to a GR value of $0$.  Nevertheless, these oscillations never end up reflecting upon the lensing potential, $\\Phi_{+}$, which remains practically indistinguishable from $\\Lambda$CDM apart from a negligible enhancement at the start of evolution. However, they do reflect upon the Newtonian potentials, translating into a signature on the ratio between them, which oscillates rather quickly, signaling a clear departure of this model from standard GR.  We also note that the evolution of $\\chi$ and $\\Phi_{+}$ in the hybrid metric-Palatini theory is related, like in metric $f(R)$ models, to the effective mass of the additional scalar degree of freedom, $m^{2}_{\\phi}$. Since the latter is smaller at early times, the range of the action of the additional fifth force will be larger. Hence, the enhancement in the perturbations, specially $\\chi$, will be greater then. And it will be greater the smaller the scale under consideration is, since these scales start their evolution deep within the range of the additional force."
    },
    "1402/1402.6341_arXiv.txt": {
        "abstract": "The far-ultraviolet (FUV; 912~--~1700~\\AA) radiation field from accreting central stars in Classical T Tauri systems influences the disk chemistry during the period of giant planet formation.  The FUV field may also play a critical role in determining the evolution of the inner disk ($r$~$<$~10 AU), from a gas- and dust-rich primordial disk to a transitional system where the optically thick warm dust distribution has been depleted. Previous efforts to measure the true stellar+accretion-generated FUV luminosity (both hot gas emission lines and continua) have been complicated by a combination of low-sensitivity and/or low-spectral resolution and did not include the contribution from the bright Ly$\\alpha$ emission line.  In this work, we present a high-resolution spectroscopic study of the FUV radiation fields of 16 T Tauri stars whose dust disks display a range of evolutionary states.  We include reconstructed Ly$\\alpha$ line profiles and remove atomic and molecular disk emission (from H$_{2}$ and CO fluorescence) to provide robust measurements of both the FUV continuum and hot gas lines (e.g., Ly$\\alpha$, \\ion{N}{5}, \\ion{C}{4}, \\ion{He}{2}) for an appreciable sample of T Tauri stars for the first time. We find that the flux of the typical Classical T Tauri Star FUV radiation field at 1 AU from the central star is $\\sim$~10$^{7}$ times the average interstellar radiation field.   The Ly$\\alpha$ emission line contributes an average of 88\\% of the total FUV flux, with the FUV continuum accounting for an average of 8\\%. Both the FUV continuum and Ly$\\alpha$ flux are strongly correlated with \\ion{C}{4} flux, suggesting that accretion processes dominate the production of both of these components. On average, only $\\sim$~0.5\\% of the total FUV flux is emitted between the Lyman limit (912~\\AA) and the H$_{2}$ (0~--~0) absorption band at 1110~\\AA.   The total and component-level high-resolution radiation fields are made publicly available in machine-readable format. ",
        "introduction": "Considerable observational effort has been invested to characterize the high-energy spectra of Classical T Tauri stars (CTTS), Class-II protostars with gas-rich circumstellar environments and active accretion.  Near-ultraviolet (NUV; $\\lambda$~=~1700~--~3200~\\AA; 7~$>$~$h\\nu$~$>$~4~eV), far-ultraviolet (FUV; $\\lambda$~=~912~--~1700~\\AA; 13.6~$>$~$h\\nu$~$>$~7~eV), extreme-ultraviolet (EUV; $\\lambda$~=~120~--~912~\\AA; 100~$>$~$h\\nu$~$>$~13.6~eV), and X-ray ($\\lambda$~$<$ 100~\\AA; $h\\nu$~$>$~0.1~keV) irradiances are important ingredients for a complete picture of the chemistry and evolution of protoplanetary environments around young stars during the epoch of giant planet formation, migration, and the growth of rocky planet cores~\\citep{ward97,armitage02,ida04}.  FUV molecular spectra provide unique insight into the physical conditions and composition of the warm inner-disk ($r$~$<$~10 AU) surface layer~\\citep{herczeg04,france11b,schindhelm12b,france12b} and molecular outflows~\\citep{walter03,herczeg06,krull07}.  NUV spectra of the Balmer continuum excess provide perhaps the most direct measure of the protostellar mass accretion rate~(e.g.; Ingleby et al. 2011b, 2013 and references therein), while detailed studies of the fluxes and profiles of FUV spectral lines~\\citep{giovannelli95,lamzin96,ardila02,gunther08,ardila13} and X-ray spectra~\\citep{schneider08,brickhouse10} provide constraints on the magnetospheric accretion paradigm.\\nocite{ingleby11b,ingleby13} Further, the intrinsic energetic radiation field, generated from a combination of accretion processes and atmospheric activity on the protostar, provide an essential input to models of the chemistry and evolution of CTTS disks.  %  {\\it Disk Chemistry~--~} The strength and shape of the FUV radiation field has a strong influence on the chemical abundances of the disk, both at planet-forming radii ($r$~$<$ 10 AU; Walsh et al. 2012) and at larger radii ($r$~$>$~50 AU) where the majority of the disk mass resides (see e.g., Bergin et al. 2007 and references therein).\\nocite{bergin07}  The stellar FUV continuum controls the dissociation of the most abundant disk molecules (H$_{2}$ and CO; Shull \\& Beckwith 1982; van Dishoeck \\& Black 1988)).  The propagation of the FUV continuum is mainly regulated by dust grains~\\citep{zadelhoff03}; the processes of grain-growth and settling likely allow these photons to penetrate deeper into the disk as the protoplanetary environment evolves~\\citep{aikawa06,vasyunin11}.  \\nocite{walsh12,shull82,vdb88}  \\citet{bergin03} first emphasized the importance of accretion-generated \\ion{H}{1} Ly$\\alpha$ to the disk chemistry, and more recently it has been shown that the FUV spectral energy distribution of $all$ CTTSs is overwhelmingly dominated ($\\gtrsim$~80\\%) by Ly$\\alpha$ emission~\\citep{schindhelm12b}.  Unlike the FUV continuum emission, the radiative transfer of Ly$\\alpha$ photons is controlled mainly by resonant scattering in the upper, atomic disk atmosphere~\\citep{bethell11}.  Subsequent detailed disk modeling has demonstrated the importance of properly accounting for Ly$\\alpha$ radiation from the central star, finding significant ($\\gtrsim$ 1 order of magnitude) depletions in the abundances of C$_{2}$H$_{4}$, CH$_{4}$, HCN, NH$_{3}$, and SO$_{2}$ when Ly$\\alpha$ is included~\\citep{fogel11}.  Interestingly, some species with large photo-absorption cross-sections at Ly$\\alpha$ ($\\lambda$~=~1216~\\AA), such as H$_{2}$O, do not show significant depletion because the enhanced dissociation rate is balanced by Ly$\\alpha$-driven photodesorption of water molecules from dust grains.  It is clear now that Ly$\\alpha$ is a mandatory component of FUV radiation fields used for chemical modeling.  However most large CTTS spectral atlases in the literature do not provide spectral coverage at 1216~\\AA\\ (e.g., Yang et al. 2012), or are dominated by geocoronal emission, such as in archival data from the {\\it International Ultraviolet Explorer}.  When Ly$\\alpha$ spectral coverage is included, scattering in the interstellar and circumstellar environment prevents a direct measurement of the local Ly$\\alpha$ environment (as with measurements from the {\\it Hubble Space Telescope}-Space Telescope Imaging Spectrograph; $HST$-STIS).\\nocite{yang12}  {\\it Disk Evolution~--~}Primordial gas disks are known to dissipate on timescales of~$\\leq$~10~Myr, at which point mass accretion onto the central star halts~\\citep{fedele10}.  While there is a growing body of evidence that inner gas disks can survive longer than the typical 2~--~4 Myr lifetime of inner dust disks (e.g., Salyk et al. 2009; France et al. 2012b)\\nocite{salyk09,france12b}, the physical process by which the inner disk is cleared is not yet established.  Various mechanisms including photoevaporation~\\citep{alexander06,gorti09} and dynamical clearing by exoplanetary systems~\\citep{rice03,dodson11}, possibly aided by a magnetorotational instability~\\citep{chiang07}, can reproduce certain transitional disk observations.~\\nocite{hernandez07,chiang07} Photoevaporation was initially considered for EUV photons from the central star~\\citep{clarke01,alexander06}, and more recent work has demonstrated that X-rays~\\citep{owen10} and FUV photons~\\citep{gorti09} can also play an important role in disk dispersal.  Models that simultaneously treat FUV, EUV, and X-ray irradiation from the central star have shown that the FUV illumination can control the total evaporation rate (and hence the disk lifetime) by driving the heating at intermediate ($r$~$\\sim$~3~--~30 AU) and large radii ($r$~$\\geq$~100 AU; Gorti \\& Hollenbach 2009).  FUV radiation also controls the gas temperature at the base of the evaporative flows through the generation of photoelectrons released by FUV-illuminated dust grains.  Grain-growth and dust settling in the disk, part of the first stages of the planet formation process, enable deeper penetration of the FUV radiation.  Therefore, planet formation itself can lead to an alteration of the temperature and chemical structure of the planet-forming region. \\citet{ingleby11} presented a comprehensive study of the evolution of FUV (excluding Ly$\\alpha$) and X-ray radiation fields over the 10 Myr lifetimes of gas disks.  However, the low-resolution data used in most previous studies suffers from molecular disk contamination (line-blending with photo-excited H$_{2}$ and CO emissions) and no previous surveys of FUV radiation fields have included a proper treatment of Ly$\\alpha$ or isolation of the FUV continuum.  Therefore, most data in the literature or readily available in the archive may not be representative of the strength $or$ the shape of the true disk-dispersing radiation fields.    {\\it This Work~--~} In order to provide a more accurate and complete (including local Ly$\\alpha$ emission profiles) observational basis for models of disk chemistry and evolution, we present new measurements of the FUV line and continuum spectra generated by accretion and magnetic processes near the protostellar photosphere.  These FUV radiation fields are available to the community in a machine-readable format\\footnote{ {\\tt http://cos.colorado.edu/$\\sim$kevinf/ctts\\_fuvfield.html} }.   These spectra were obtained with the Cosmic Origins Spectrograph (COS) and the STIS aboard $HST$.  While {\\it Far-Ultraviolet Spectroscopic Explorer} ($FUSE$) observations of a small number of bright CTTSs are available in the literature~\\citep{wilkinson02,herczeg04,herczeg05,herczeg06,gunther08}, $FUSE$ did not have the sensitivity to study ``typical'' Taurus-Auriga CTTSs in detail.  We therefore use $FUSE$ observations to constrain our 912~--~1150~\\AA\\ radiation field creation (in particular the shapes and strengths hot gas emission lines and continua), but we present an approach where the short-wavelength radiation fields are inferred from longer-wavelength $HST$ data.  \\begin{figure} \\figurenum{1} \\begin{center} \\epsfig{figure=f1.eps,width=3.4in,angle=00} \\vspace{+0.25in} \\caption{Complete FUV spectra of the 16 CTTSs studied in this work, including reconstructed Ly$\\alpha$ emission lines. These spectra are coadditions of $HST$-COS observations in the G130M and G160M modes (except for TW Hya, which was observed with STIS E140M; Herczeg et al. 2002) at several central wavelengths and focal-plane split positions.  Almost all of the structures seen in these data are real atomic and molecular emission and absorption features. The data have been corrected for interstellar reddening (Table 1), scaled to the flux at 1 AU from the central star for comparison, and binned by three spectral resolution elements (21 pixels) for display. \\label{cosovly}} \\end{center} \\end{figure}   \\begin{figure} \\figurenum{2} \\begin{center} \\epsfig{figure=f2.eps,width=3.4in,angle=00} \\vspace{+0.25in} \\caption{The binned FUV continuum spectra are shown as gray filled circles.  A second order polynomial fit is extrapolated down to the Lyman Limit (912~\\AA) and is shown as the red dashed line.   The ``1600~\\AA\\ Bump'' (spanning $\\sim$~1520~--~1660~\\AA) is prominent (detected at $>$ 3$\\sigma$ significance) in  10/16 targets. \\label{cosovly}} \\end{center} \\end{figure}  In Section 2, we describe the targets and the $HST$ observations.  In Section 3, we briefly describe the spectral deconvolution performed to separate the continuum, hot gas atomic line, and molecular line emission.   The details of this deconvolution can be found in the Appendix.  We detect the quasi-continuous emission feature near 1600~\\AA, the ``1600~\\AA\\ Bump'' in $\\sim$~70\\% of the sources; the analysis of this emission~\\citep{bergin04,ingleby09,france11a}  will be the subject of a future work.  In Section 4, we discuss the stellar+accretion spectra and present correlations suggesting that the majority of the Ly$\\alpha$ and FUV continuum are generated by accretion processes.  % We present a summary of this work in Section 5. ",
        "conclusions": "There are numerous indicators of the gas and dust content of a young protoplanetary system. Three important observables are the warm dust content of the inner disk, the presence of circumstellar gas, and signs of active accretion. Our $HST$ observations provide measurements of the last two, while the first has been extensively studied in the IR. % The large transition probabilities of the H$_{2}$ electronic band systems and the lack of photospheric emission at $\\lambda$~$<$~1700~\\AA\\ in low-mass stars make fluorescent H$_{2}$ one of the most sensitive indicators for the presence of molecular gas in the inner ~$\\sim$~10~AU of young circumstellar disks~\\citep{france12b}. H$_{2}$ emission line spectroscopy can directly probe gas surface densities as small as $\\Sigma_{H2}$~$\\lesssim$~10$^{-6}$ g cm$^{-2}$. The FUV bandpass also contains resonance lines of hydrogen, carbon, nitrogen, and oxygen which are powerful diagnostics of warm/hot gas ($T_{form}$~$\\sim$~10$^{4}$~--~3~$\\times$~10$^{5}$ K; \\ion{H}{1} Ly$\\alpha$, \\ion{C}{3} $\\lambda$977, \\ion{C}{4} $\\lambda$$\\lambda$1548,1550, \\ion{N}{5} $\\lambda$$\\lambda$1239,1243, and \\ion{O}{6} $\\lambda$$\\lambda$1032,1038; as well as the H$\\alpha$ line of ionized helium, \\ion{He}{2} $\\lambda$1640) formed near the accretion shock region\\footnote{including the heated photosphere at the base of the accretion column, the pre-shock region, shock surface, and post-shock region}  by collisional- and photo- excitation and ionization processes. Our observations enable the first measurements of the FUV continuum in many of these targets, a potentially useful diagnostic of the accretion environment that has only recently been made accessible by the high sensitivity and spectral resolution of $HST$-COS (e.g., France et al. 2011a).\\nocite{france11b}   \\subsection{Ly$\\alpha$, Hot Gas Lines, and the FUV Continuum: Constraints on Physical Origins % }  By bringing together the reconstructed Ly$\\alpha$ emission profiles, the ``H$_{2}$ cleaned'' hot gas emission lines, and measurements of the uncontaminated FUV continuum, we can explore the relations between these different components and empirically constrain the origin of the strong Ly$\\alpha$ and FUV continuum emission in these sources.  The prevailing picture for the formation of the UV continuum excess is that it arises from ionized gas ($T_{cont}$~$\\sim$~1.5~--~3~$\\times$~10$^{4}$ K) in an optically thin pre-shock region at the base of the accretion column~\\citep{calvet98,costa00}.  In this picture, the FUV continuum may be an extension of the Balmer continuum used for accretion rate determination in the NUV (see e.g., Ingleby et al. 2013); and the connection between the NUV and FUV continua has been explored previously for DF Tau~\\citep{france11a}.   In the following subsection, we show that for a subsample of six of our targets with nearly simultaneous FUV and NUV observations, the FUV and NUV continua are characterized by a significantly different slope.  This demonstrates that the connection proposed by~\\citet{france11a} was most likely an artifact of comparing non-simultaneous observations, although this cannot be conclusively demonstrated for all of our targets in our sample without contemporaneous FUV and NUV data.  We find that while some of our sources display a monotonically decreasing FUV continuum towards shorter wavelengths (BP Tau, DF Tau, RU Lup, see also Herczeg et al. 2005), many of the FUV continua are flat across the FUV band (GM Aur, RECX-15, UX Tau, V4046 Sgr).  This suggests that multiple physical processes likely contribute in the complex emitting region near the stellar photosphere and the inner accretion disk.  Our aim here is to provide the data for use in detailed models of disk chemistry and evolution, and therefore we do not attempt to present a comprehensive model for the continuum emission region.  The hot gas emissions from \\ion{C}{4} and \\ion{N}{5} have been long considered to be related to accretion processes~\\citep{krull00}. A detailed spectral line analysis of a larger sample of COS and STIS observations has shown that these lines are likely produced by hot gas in accretion spots near the stellar photosphere, the edges of accretion columns, or multiple accretion columns of varying densities~\\citep{ardila13}.  The \\ion{He}{2} lines however appear to be dominated by a combination of emission from the magnetically active stellar atmosphere and the pre-shock region.  Ly$\\alpha$ is presumably generated at several places in the near-star environment, however we are not aware of any previous work that explores the observational connection between Ly$\\alpha$ luminosity and accretion processes (see Muzerolle et al. 2001 and Kurosawa et al. 2006 for a detailed description of Balmer lines).\\nocite{muzerolle01,kurosawa06} \\begin{figure*} \\figurenum{6} \\begin{center} \\epsfig{figure=f6.eps,width=5.0in,angle=90} \\vspace{+0.25in} \\caption{ A comparison between the reddening corrected FUV (gray filled circles) and NUV (black line) continua for the six stars in our sample with contemporaneous observations.  The solid orange line and dashed pink lines are linear fits to the FUV and NUV continua, respectively.  The linear fit to the FUV data is created using fluxes from 1150~--~1174~\\AA\\ and 1680~--~1750~\\AA, excluding contributions from the molecular continuum and the ``1600~\\AA\\ Bump''.  The linear fits show that the FUV continuum is brighter than suggested by an extrapolation of the NUV continuum for most targets, and that the FUV slopes are shallower, possibly suggesting multiple emitting regions. \\label{cosovly}} \\end{center} \\end{figure*}        In order to empirically constrain the relationship between Ly$\\alpha$ and the FUV continuum and accretion processes in the following subsections, we assume that the integrated \\ion{C}{4} flux is a proxy for the mass accretion rate~\\citep{krull00,ardila13}.  It is critical to have this quasi-simultaneous accretion diagnostic because accretion rates and UV fluxes from CTTSs are known to vary by factors of $\\sim$~2 on timescales of months-to-years~\\citep{valenti00}.  Therefore, accretion rates in the literature will not necessarily be well correlated with the UV emission at the time of our observations.  In Figure 5 {\\it top left}, we demonstrate the tight correlation between the accretion-dominated \\ion{C}{4} and \\ion{N}{5} flux for our stars (see also the larger sample from Ardila et al. 2013).  For all of the correlations shown in Figure 5, we present the Spearman rank correlation coefficient ($\\rho$) and the likelihood that this population is drawn from a random sample (i.e., not correlated; $n$).  In general,  $\\vert$$\\rho$$\\vert$~$>$~0.5 and $n$~$<$~0.05 indicates that a real correlation exists between the two quantities.   For instance, [$\\rho$,$n$] = [0.81, 1.3~$\\times$~10$^{-4}$] for the plot of \\ion{N}{5} and \\ion{C}{4}, indicating a strong correlation, as expected.   \\subsubsection{FUV Continuum}  The FUV continuum flux (in units of erg cm$^{-2}$ s$^{-1}$), evaluated at 1 AU from the central star, is computed by integrating the polynomial fit to the binned intra-emission-line $HST$ spectra (\\S A.3) from 912~--~1650~\\AA\\ (Figures 2 and A.4~--~A.6).  The FUV continuum is shown to be tightly correlated with the \\ion{C}{4} emission (Figure 5, {\\it top right}, [$\\rho$,$n$] = [0.93, 1.5~$\\times$~10$^{-7}$]).  We interpret this as strong evidence that the CTTS FUV continuum is generated by accretion processes, either directly or powered by accretion luminosity.  However, it may originate in a spatially separate region from the NUV Balmer continuum, as mentioned above.  In Figure 6, we compare data for six of our targets (DM Tau, DR Tau, GM Aur, HN Tau, RECX-11, and RECX-15) with FUV and NUV observations separated by less than a few hours.  Linear fits to the FUV (excluding the molecular quasi-continuum) and NUV regimes are shown overplotted as solid orange and dashed pink lines, respectively.  The observed FUV continuum is shown to be brighter than predicted from a simple extrapolation of the NUV continuum, and has a significantly shallower slope (see also Herczeg et al. 2004, 2005).\\nocite{herczeg04,herczeg05} We note that the S/N ratios of individual binned data points are relatively high, demonstrating that the excess FUV continuum flux is statistically significant in all cases. The average S/N per binned continuum point ranges from 2.5~--~14.8 from 1140~--~1340~\\AA\\ and 0.9~--~17.0 from 1660~--~1740~\\AA\\ (Table A.1). While it seems clear that the FUV continuum is related to accretion processes, these differences argue that the FUV and NUV continua may be formed in spatially distinct regions near the interaction of the accretion streams and the stellar photosphere.    Recent models of the NUV Balmer excess have incorporated multiple accretion components with a range of energy fluxes and densities (e.g., Ingleby et al. 2013), and future work may be able to connect the FUV and NUV continua by incorporating these different components.  Our analysis finds that the FUV continuum extends to $\\lambda$~$<$~1216~\\AA\\ in all cases, ruling out the possibility that this flux originates from \\ion{H}{1} 2-photon (2$s$~$^{2}S_{1/2}$~--~1$s$~$^{2}S_{1/2}$) emission.  Furthermore, \\ion{H}{1} 2-photon emission is suppressed in high-density regions ($n_{H}$~$>$~1.5~$\\times$~10$^{4}$ cm$^{-3}$), consistent with the scenario in which this emission is the high-frequency extension of the Balmer continuum generated in the relatively dense pre-shock region at the base of the accretion column~\\citep{calvet98}.  A future space-borne instrument capable of simultaneous FUV and NUV spectroscopy or photometry could provide more insight on the potential link between the FUV continuum and the NUV Balmer continuum in CTTSs.   \\begin{deluxetable*}{lccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Relative Contributions to the Stellar + Accretion FUV Radiation Field\\label{lya_lines}} \\tablewidth{0pt} \\tablehead{ \\colhead{Target} & \\colhead{$F_{tot}$\\tablenotemark{a}} &   \\colhead{log$_{10}$($F_{tot}$/G$_{o}$)\\tablenotemark{b} }   & \\colhead{FUV Continuum} & \\colhead{Ly$\\alpha$ \\tablenotemark{c}} & \\colhead{\\ion{C}{4}} & \\colhead{Other lines\\tablenotemark{d}}  & \\colhead{$\\lambda$~$\\leq$~1110~\\AA}  \\\\ &    (erg cm$^{-2}$ s$^{-1}$)    &      &    (\\%)    &     (\\%)     &    (\\%)     &      (\\% )  &      (\\% )   } \\startdata AATAU & 3.1~$\\times$~10$^{4}$   &   7.3   &   2.8   &   95.9   &   0.8   &   0.6  &  0.1   \\\\ BPTAU & 3.5~$\\times$~10$^{4}$   &   7.3   &   16.8   &   75.7   &   4.5   &   3.1 &  1.8  \\\\ DETAU & 1.5~$\\times$~10$^{4}$   &   7.0   &   9.2   &   88.0   &   2.0   &   0.9  &  0.3 \\\\ DFTAU & 2.3~$\\times$~10$^{5}$   &   8.2   &   2.1   &   97.2   &   0.4   &   0.2   &  0.1\\\\ DMTAU & 4.3~$\\times$~10$^{3}$   &   6.4   &   8.4   &   88.0   &   1.7   &   1.9 &  0.3  \\\\ DRTAU & 4.0~$\\times$~10$^{4}$   &   7.4   &   46.2   &   49.4   &   3.3   &   1.0  &  0.7 \\\\ GMAUR & 1.2~$\\times$~10$^{4}$   &   6.9   &   9.9   &   86.2   &   2.4   &   1.5  &  0.4 \\\\ HNTAU & 1.2~$\\times$~10$^{4}$   &   6.9   &   10.1   &   88.4   &   0.9   &   0.5 &  0.1  \\\\ LKCA15 & 1.6~$\\times$~10$^{4}$   &   7.0   &   6.5   &   90.4   &   2.1   &   1.0  &  0.3  \\\\ RECX11 & 2.3~$\\times$~10$^{3}$   &   6.2   &   4.6   &   90.9   &   2.9   &   1.7  &  0.6 \\\\ RECX15 & 4.8~$\\times$~10$^{3}$   &   6.5   &   2.8   &   96.7   &   0.3   &   0.3  &  0.1 \\\\ RULUPI & 2.8~$\\times$~10$^{4}$   &   7.2   &   17.6   &   79.8   &   2.0   &   0.7  &  0.3 \\\\ SUAUR & 3.3~$\\times$~10$^{4}$   &   7.3   &   8.4   &   88.3   &   2.1   &   1.2  &  0.4 \\\\ TWHYA & 9.7~$\\times$~10$^{3}$   &   6.8   &   16.7   &   71.9   &   6.5   &   4.9  &  0.9 \\\\ UXTAU & 5.6~$\\times$~10$^{3}$   &   6.5   &   4.8   &   92.2   &   1.7   &   1.3   &  0.3\\\\ V4046SGR & 1.5~$\\times$~10$^{4}$   &   7.0   &   5.3   &   91.8   &   1.2   &   1.8 &  0.9 \\\\ \\tableline Average\\tablenotemark{e}  &     &  7.0~$\\pm$~0.5    &   8.4~$\\pm$~5.2   &   88.1~$\\pm$~7.3   &   2.1$~\\pm$~1.6   &  1.4~$\\pm$~1.2  &  0.5~$\\pm$~0.4 \\enddata \\tablenotetext{a}{Integrated 912~--~1650~\\AA\\ stellar+accretion FUV radiation field (not including molecular fluorescence lines, low-ionization atomic emission, or the ``1600~\\AA\\ Bump''), evaluated at 1 AU from the central pre-main sequence star.  } \\tablenotetext{b}{Ratio of the integrated FUV radiation field at 1 AU to the average interstellar radiation field (1.6~$\\times$~10$^{-3}$ erg cm$^{-2}$ s$^{-1}$; Habing 1968). } \\tablenotetext{c}{Intrinsic Ly$\\alpha$ emission, integrated over 1211~--~1221~\\AA .} \\tablenotetext{d}{Other stellar+accretion hot gas emission lines, \\ion{C}{3} $\\lambda$977~\\AA\\ + \\ion{O}{6} $\\lambda$$\\lambda$1032,1038~\\AA\\ + \\ion{N}{5} $\\lambda$$\\lambda$1239, 1243~\\AA\\ + \\ion{He}{2} $\\lambda$1640~\\AA. } \\tablenotetext{e}{Average quantities are calculated excluding DR Tau, whose Ly$\\alpha$ profile reconstruction is compromised by multiple Ly$\\alpha$ emission sources (\\S A.1).} \\end{deluxetable*}   \\subsubsection{Ly$\\alpha$ Line Emission}  The intrinsic Ly$\\alpha$ flux (in units of erg cm$^{-2}$ s$^{-1}$), evaluated at 1 AU from the central star, is computed by integrating the reconstructed Ly$\\alpha$ line profile from 1211~--~1221~\\AA.  We note that we take the intrinsic reconstructed line profile as opposed to the line profile seen by the H$_{2}$ molecules at the disk surface after absorption by atomic hydrogen in the outflow (see Figure A.1).  A discussion of the Ly$\\alpha$ profiles including the neutral outflow absorption component is presented by~\\citet{schindhelm12b}, and the uncertainty on the intrinsic Ly$\\alpha$ flux used here is $\\sim$~10~--~30\\%, depending on the detectability and S/N of the H$_{2}$ fluorescence lines used for the reconstruction~\\citep{france12b}.  The middle panels of Figure 5 show the correlation between Ly$\\alpha$ and the FUV continuum ([$\\rho$,$n$] = [0.71, 1.9~$\\times$~10$^{-3}$]) and the \\ion{C}{4} emission ([$\\rho$,$n$] = [0.77, 4.8~$\\times$~10$^{-4}$]).  The Ly$\\alpha$ flux is observed to correlate with both quantities, but with less significance than the correlations between \\ion{C}{4} and \\ion{N}{5}, and \\ion{C}{4} and the FUV continuum.  This is possibly confirmed by another representation on the bottom panels of Figure 5, where the ratio of Ly$\\alpha$/FUV continuum is plotted against the published mass accretion rates ($left$, [$\\rho$,$n$] = [$-$0.31, 0.24]) and the \\ion{C}{4} fluxes ($right$, [$\\rho$,$n$] = [$-$0.50, 0.047]).  The plots are qualitatively similar; Figure 5 ({\\it bottom left}) does not show a significant anti-correlation between Ly$\\alpha$/FUV continuum and the mass-accretion rate, Ly$\\alpha$/FUV continuum and \\ion{C}{4} show a weak anti-correlation suggesting that the FUV continuum is more tightly correlated with the accretion than the Ly$\\alpha$ emission.  \\begin{figure} \\figurenum{7} \\begin{center} \\epsfig{figure=f7.eps,width=3.5in,angle=00} \\vspace{+0.0in} \\caption{ A comparison between observed CTTS and WTTS Ly$\\alpha$ profiles for a few example objects.  The CTTS Ly$\\alpha$ profiles (GM Aur and BP Tau) uniformly display higher fluxes and line widths ($\\Delta$$v$~$>$~$\\pm$~500 km s$^{-1}$) than their WTTS counterparts (TWA 7 and LkCa19), suggesting that accretion processes contribute (and probably dominate) to the total Ly$\\alpha$ luminosity. The narrow central emission is from Earth's geocorona and the broad depression seen at the core of the CTTS line profiles is due to resonant scattering in the interstellar medium along the line-of-sight. \\label{cosovly}} \\end{center} \\end{figure} The production of Ly$\\alpha$ is still likely dominated by accretion processes.  Figure 7 shows the observed Ly$\\alpha$ emission lines from two representative CTTSs from this work (GM Aur and BP Tau) and two WTTSs (TWA 7 and LkCa19).  In all cases, the CTTS Ly$\\alpha$ lines are much stronger and broader than WTTS Ly$\\alpha$ profiles for a similar age and stellar mass.  All CTTSs display broad Ly$\\alpha$ wings to $v_{Ly\\alpha}$~$>$~500 km s$^{-1}$, and many show Ly$\\alpha$ emission wings in excess of $\\pm$~1000 km s$^{-1}$. There may be several physical reasons that the correlation between Ly$\\alpha$ and \\ion{C}{4} is weaker than between the FUV continuum and \\ion{C}{4}, including uncertainties in the Ly$\\alpha$ reconstruction process, particularly when there are multiple sources of Ly$\\alpha$ photons in the systems (stellar atmosphere, outflows and/or disk surface) and resonant scattering of Ly$\\alpha$ photons in the gas-rich CTTS circumstellar environment.  \\subsection{Absolute Fluxes at Planet-forming Radii and Relative Contributions}  In order to understand the UV-driven photochemistry of the protoplanetary environment, one needs constraints on both the absolute flux and the shape of the illuminating radiation field.  The absolute flux of the radiation environment determines the photoexcitation and photodissociation rates for disk molecules and the SED determines the primary scattering and absorption mechanisms for diffusing FUV photons through the disk~\\citep{bergin03,bethell11,fogel11}.  In Table 2, we give the total, reddening-corrected flux from our targets (evaluated at 1 AU to enable inter-comparison) integrated over the 912~--~1650~\\AA\\ bandpass.  In the second column, this quantity is given in units of the average interstellar radiation field from~\\citet{habing68}.  The average incident flux from the accreting central star is~$\\sim$~10$^{7}$ times the contribution from scattered OB starlight at 1 AU from the central star, although the shapes of the CTTS radiation field and the interstellar field are very different.  By spectrally resolving the various components of the FUV field and including the reconstructed Ly$\\alpha$ emission, we can provide an inventory of the FUV luminosity sources in $\\sim$~1~--~10 Myr CTTSs.  Table 2 shows the relative contribution of each component to the total FUV irradiance from each target.  As argued by previous authors~\\citep{bergin03,herczeg04,schindhelm12b}, Ly$\\alpha$ dominates the FUV radiation output from these sources, with an average fractional luminosity of 88.1~$\\pm$~7.3~\\%\\footnote{The only object for which Ly$\\alpha$ may not be the largest contributor to the FUV radiation field is DR Tau. However, the Ly$\\alpha$ reconstruction for this star is rather uncertain due to the presence of multiple sources of Ly$\\alpha$ emission in the system, most likely a second, narrow component associated with an outflow. }. The FUV continuum is the second largest energy source, with 8.4~$\\pm$~5.2~\\%.  The contribution from the \\ion{C}{4} $\\lambda$$\\lambda$1548,1550~\\AA\\ doublet is 2.1~$\\pm$~1.6~\\%, and the remaining hot gas emission lines make up $<$~5\\% in all cases, although the exact contribution is rather uncertain because we assumed the TW Hya \\ion{C}{3} and \\ion{O}{6} fluxes, scaled by the relative \\ion{C}{4} fluxes, for the remaining targets.  Lower ionization species likely contribute $<$ 1\\% to the total field strength (see A.2).  The primary transitions leading to dissociation of H$_{2}$ and CO, through the H$_{2}$ $B$~--~$X$ and $C$~--~$X$ band systems (radiatively dissociative) and the CO $E$~--~$X$ band system (predissociative), are located at $\\lambda$~$\\leq$~1110~\\AA.  The $\\lambda$~$\\leq$~1110~\\AA\\ radiation field also plays an important role in the production of photo-electrons that heat the gas in the disk surface layers~\\citep{pedersen11}.  The 912~--~1110~\\AA\\ bandpass contains between 0.1~\\% and 1.8~\\% of the total FUV flux, with an average (excluding DR Tau) of 0.5~$\\pm$~0.4~\\%.   On average, the 912~--~1110~\\AA\\ CTTS radiation field has a flux of $\\sim$~150 erg cm$^{-2}$ s$^{-1}$ at 1 AU from the central star.     \\subsection{Comparison with Low-resolution Spectra}  In this subsection, we quantify the impact of low-resolution data on the determination of the total FUV flux from CTTSs.  Previous large FUV spectroscopic samples of CTTSs were acquired at low spectral resolution ($R$~$\\sim$~80~--~1000; Johns-Krull et al. 2000; Valenti et al. 2000; Ingleby et al. 2009, 2011; Yang et al. 2012).  Most of these observations did not have access to Ly$\\alpha$ directly due to instrumental bandpass limitations and/or contamination by geocoronal emission.  Furthermore, low-resolution prevents the necessary measurements of individual H$_{2}$ or CO fluorescent emission lines necessary for reconstructing the Ly$\\alpha$ emission profile.  Consequently, most quoted measurements of the FUV radiation field strength are integrated over $\\approx$ 1230~--~1700~\\AA~\\citep{ingleby11,yang12}, $excluding$ the dominant FUV flux source, \\ion{H}{1} Ly$\\alpha$.  Low resolution data are also insufficient for separating the accretion-powered FUV continuum from the wealth of narrow molecular emission lines and the molecular continuum (the ``1600~\\AA\\ Bump''). In order to quantify the differences between high- and low-resolution CTTS FUV radiation fields, data for 10 objects with previous low-resolution FUV observations from $HST$ were retrieved from the~\\citet{yang12} atlas of low-resolution T Tauri star spectra hosted at MAST.  The average flux ratio between the observed low-resolution 1235~--~1700~\\AA\\ flux from the~\\citet{yang12} atlas ($F(1235-1700)_{LR}$) and the 1235~--~1700~\\AA\\ fluxes from this work ($F(1235-1700)_{fit}$), $F(1235-1700)_{LR}$/$F(1235-1700)_{fit}$, is 1.41~$\\pm$~0.52 (Table 3).  The average being greater than unity most likely reflects the inclusion of Ly$\\alpha$-pumped H$_{2}$ and CO fluorescence in low-resolution data and the scatter is most likely the time-variability that is characteristic of these sources (recall these data were acquired over a multi-year baseline; and see e.g., Giovannelli et al. 1995; G{\\'o}mez de Castro \\& Fern{\\'a}ndez 1996).  The two sources which display the largest FUV variability are DM Tau and SU Aur, showing a factor of $\\sim$~3 variability between the two epochs.  If we screen sources farther than 3--$\\sigma$ from the mean of the distribution,  the average drops to $F(1235-1700)_{LR}$/$F(1235-1700)_{fit}$ = 1.06~$\\pm$~0.34.  \\begin{deluxetable}{lccc} \\tabletypesize{\\normalsize} \\tablecaption{Comparison to Low-Resolution CTTS Spectra \\label{lya_lines}} \\tablewidth{0pt} \\tablehead{ \\colhead{Target} & \\colhead{$HST$ Mode\\tablenotemark{a}} & \\colhead{ $\\frac{ F(1235-1700)_{LR} }{ F(1235-1700)_{fit} }$ \\tablenotemark{b} }   & \\colhead{ $\\frac{ F(1235-1700)_{LR} }{ F(1150-1700)_{fit} }$ \\tablenotemark{c} } } \\startdata AATAU & ACS   &   1.18   &   0.07   \\\\ BPTAU & STIS   &   0.70   &   0.21     \\\\ DETAU & ACS   &   1.68   &   0.27   \\\\ DMTAU & STIS   &   2.73   &   0.31   \\\\ DRTAU & ACS   &   0.77   &   0.55    \\\\ GMAUR & ACS   &   0.99   &   0.14     \\\\ GMAUR & STIS   &   0.96  &   0.14     \\\\ HNTAU & ACS   &   1.04   &   0.17    \\\\ LKCA15 & STIS   &   1.43   &   0.19     \\\\ SUAUR & STIS  &   3.24   &   0.60   \\\\ V4046SGR & ACS   &   0.76   &   0.05   \\\\ \\tableline Average   &     &  1.41~$\\pm$~0.52    &   0.25~$\\pm$~0.17 \\enddata \\tablenotetext{a}{$HST$ Low-resolution mode used for comparison. Low-resolution data taken from~\\citet{yang12}.   ACS refers to ACS/SBC PR130L mode, STIS refers to STIS G140L. } \\tablenotetext{b}{Ratio of the observed low-resolution spectrum from 1235~--~1700~\\AA\\ to the extracted, intrinsic CTTS spectrum from 1235~--~1700~\\AA. } \\tablenotetext{c}{Ratio of the observed low-resolution spectrum from 1235~--~1700~\\AA\\ to the extracted, intrinsic CTTS spectrum from 1150~--~1700~\\AA, including the reconstructed Ly$\\alpha$ emission line.  % } \\end{deluxetable}  Comparing the ratio of the observed low-resolution 1235~--~1700~\\AA\\ flux from the~\\citet{yang12} atlas and our 1150~--~1700~\\AA\\ fluxes that {\\it include the reconstructed Ly$\\alpha$ profiles} ($F(1150-1700)_{fit}$), we find  $F(1235-1700)_{LR}$/$F(1150-1700)_{fit}$~=~0.25~$\\pm$~0.17.  Applying the 3--$\\sigma$ screen, $F(1235-1700)_{LR}$/$F(1150-1700)_{fit}$~=~0.17~$\\pm$~0.08.  The general result is that even with the additional contribution from H$_{2}$ emission, and the uncertainty associated with the time-variability of the sources, low-resolution FUV spectra that do not include a reconstructed Ly$\\alpha$ emission line underestimate the total FUV flux by a factor of $\\approx$~6.  While a factor of $\\approx$~6 is the average flux underestimate, Table 3 shows that in the extreme cases (e.g., AA Tau and V4046 Sgr), the low-resolution 1235~--~1700~\\AA\\ data underestimate the true intrinsic stellar+accretion FUV radiation field strength by factors of ~$\\geq$~15.  Our conclusion from this comparison with low-resolution data in the literature is that ACS/SBC PR130L and STIS G140L-based flux estimates underestimate the true FUV radiation field strength from CTTSs by approximately an order-of-magnitude.  Our findings quantitatively support the assertions of previous authors~\\citep{bergin03,herczeg04,schindhelm12b} that Ly$\\alpha$ must be included to make an accurate of estimate of $both$ the strength and the shape of the FUV radiation field in protoplanetary environments around low-mass stars.  It should also be emphasized that the largest uncertainty on the inferred local FUV radiation field strength is likely the uncertainty on the line-of-sight reddening, both interstellar and circumstellar.  Typical dispersions on the derived values of $A_{V}$ in the literature are 0.3~--~1.0 magnitudes, the relative contributions of interstellar and circumstellar grains are unclear~(see e.g., McJunkin et al. 2014 and references therein), and the scattering and extinction properties of grains in regions like Taurus may be quite different than those found in the diffuse interstellar medium~\\citep{calvet04}.\\nocite{mcjunkin14}  With empirically-derived UV radiation fields now available, the uncertainty on the extinction is likely the limiting factor on our knowledge of the absolute flux of the energetic radiation environments around young stars."
    },
    "1402/1402.6031_arXiv.txt": {
        "abstract": "We perform $N$-body simulations on a multiple massive black hole (MBH) system in a host galaxy to derive the criteria for successive MBH merger. The calculations incorporate the dynamical friction by stars and general relativistic effects as pericentre shift and gravitational wave recoil.  The orbits of MBHs are pursed down to ten Schwarzschild radii ($\\sim 1$AU).  As a result, it is shown that about a half of MBHs merge during 1~Gyr in a galaxy with mass $10^{11}M_{\\odot}$ and stellar velocity dispersion $240$~km~s$^{-1}$, even if the recoil velocity is two times as high as the stellar velocity dispersion. The dynamical friction allows a binary MBH to interact frequently with other MBHs, and then the decay of the binary orbits leads to the merger through gravitational wave radiation, as shown by \\citet{Tanikawa11}. We derive the MBH merger criteria for the masses, sizes, and luminosities of host galaxies. It is found that the successive MBH mergers are expected in bright galaxies, depending on redshifts.  Furthermore, we find that the central stellar density is reduced by the sling-shot mechanism and that high-velocity stars with $\\sim 1000$~km~s$^{-1}$ are generated intermittently in extremely radial orbits. ",
        "introduction": "Massive black holes (hereafter MBHs) with the mass of more than $10^6$ solar mass ($M_{\\odot}$) have been found in the centres of galaxies. The mass of MBHs is correlated with the properties of the spheroidal components of their host galaxies, with respect to the mass \\citep{Kormendy95,Magorrian98,Marconi03}, the velocity dispersions \\citep{Ferrarese00,Tremaine02,Gultekin09}, and the number of globular clusters \\citep{Burkert10,Harris11}. The origin of MBHs is an open issue of great significance.  In the last decade, quasars (QSOs) that possess $\\sim 10^9~M_{\\odot}$ MBHs have been found at high redshifts of $z \\gtrsim 6$ \\citep[e.g.,][]{Fan01}, that is, at the cosmic age of $\\lesssim 1$~Gyr. Conservatively speaking, the seeds of the MBHs could be stellar mass black holes as massive star remnants. In particular, the remnants of first stars are one of plausible candidates, since first stars are likely to be as massive as a few hundred solar mass \\citep{Abel00,Nakamura01,Bromm02,Yoshida06}, several tens solar mass \\citep{Clark11}, or about $50M_{\\odot}$ \\citep{Hosokawa11}. However, in order for first star remnants to grow up to $\\sim 10^9~M_{\\odot}$ in $1$~Gyr, the Eddington ratio of mass accretion rate should be larger than unity \\citep[e.g.,][and references therein]{Umemura01,Greene12}. Super-Eddington accretion is one of possible solutions for the MBH growth \\citep[e.g.][]{Abramowicz88,Kawaguchi03,Ohsuga05}.  On the other hand, the integration of the QSO luminosity function is concordant with the integrated mass function of MBHs in the local universe, as long as the Eddington ratios are between 0.1 and 1.7 \\citep{Soltan82,Yu02,Marconi04}. This implies that supermassive black holes acquire the bulk of mass through gas accretion in the late evolutionary stages and the mass accretion rates are not highly super-Eddington. Also, the gas accretion onto the seeds should be intermittent, and on average could be lower than the Eddington accretion rate \\citep{Milosavljevic09a,Milosavljevic09b}.  If the merger of multiple black holes precedes the growth via gas accretion, the merged MBH can be a seed of a supermassive black hole, and therefore the constraint for the BH growth can be weaker.  In the cold dark matter cosmology, larger galaxies form hierarchically through mergers of smaller galaxies. Hence, many MBHs are assembled in a larger galaxy, if smaller galaxies already possess MBHs. Furthermore, MBHs could be born in hyper-massive star clusters formed by galaxy collisions \\citep{Matsui11}. Thus, galaxy merger remnant can contain many MBHs, even if precursory galaxies have no MBHs.  Observationally, multiple active galactic nucleus (AGN) systems have been discovered recently.  They include a triple AGN in the galaxy SDSS J1027+1749 at $z = 0.066$ \\citep{Liu11}, three rapidly growing MBHs of $10^6-10^7 M_\\odot$ in a clumpy galaxy at $z = 1.35$ \\citep{Schawinski11}, a firstly-discovered physical quasar triplet QQQ J1432--0106 within the projected separation of $30-50$kpc at $z = 2.076$ \\citep{Djorgovski07}, and a second quasar triplet QQQ J1519+0627 within the projected separation of $200$kpc, which is likely to be harboured in a yet-to-be-formed massive system at $z=1.51$ \\citep{Farina13}.  According to the hierarchical merger history, galaxies with many MBHs are likely to form at higher redshifts. Although the galaxy merger proceeds through the violent relaxation, the merger of MBHs has difficulty.  As pointed out in \\citet{Begelman80}, two MBHs in a galaxy are likely to form a binary, but unlikely to merge directly due to the so-called loss cone depletion (the depletion of stars on orbits that intersect the binary MBH).  A binary MBH cannot reach sub-parsec separation due to the loss cone depletion, which is called the final parsec problem \\citep[e.g.][]{Merritt04}.  A possible way to evade the loss cone depletion is the nonaxisymmetric potential of the host galaxy \\citep{Merritt04,Berczik06,Khan11,Khan12}, which is the natural consequence of the galaxy merger.  A binary MBHs can merge in the nonaxisymmetric potential in $10$~Gyr or $0.3$~Gyr, when the galaxy contains stars with $10^9M_{\\odot}$ or $10^{11}M_{\\odot}$, respectively \\citep{Khan11}.  However, since this timescale is comparable to or longer than the galactic dynamical timescale, other galaxies harbouring MBHs can intrude before two MBHs merge.  This is likely to occur at higher redshifts of $z \\gtrsim 6$, at which the universe age is less than $1$~Gyr. If there are more than two MBHs in a galaxy, the dynamical relaxation of MBHs is significantly controlled by the gravity of MBHs themselves, especially by three-body interaction.  When a third MBH intrudes into a binary MBH, one of the three MBHs carries away angular momentum from the rest two MBHs, reducing the binary separation, and eventually the binary merges \\citep[e.g.][]{Iwasawa06}.  So far, galaxy structures have not been investigated when the galaxies contain more than three MBHs, although they have been investigated in the cases of two MBHs \\citep[e.g.][]{Khan11,Khan12} and three MBHs \\citep[e.g.][]{Iwasawa08}.  \\cite{Tanikawa11} (hereafter, paper~I) scrutinised a system of multiple MBHs in a galaxy by high-resolution $N$-body simulations, and found that multiple MBHs produce one dominant MBH through successive mergers. Binary MBHs lose their angular momentum owing to sling-shot mechanism, which induces the decay of the binary orbits through gravitational wave (GW) radiation. In paper~I, we investigated one model of a galaxy containing multiple MBHs. In this paper, we explore the evolution of multiple MBHs in galaxies with different 3-dimensional stellar velocity dispersions to derive the criteria of the MBH merger. We also consider the effect of the recoil by anisotropic GW radiation at the MBH merger.  Since the recoil velocity typically reaches several hundred km~s$^{-1}$ \\citep{Kesden10}, it could suppress the MBH growth. Furthermore, we investigate the impact by the MBH merger on the galaxy structure.  The paper is organised as follows. In section \\ref{sec:model}, we describe the simulation model. In section \\ref{sec:results}, we show numerical results. In section \\ref{sec:criteria}, the results are translated to derive the criteria for MBH merger, which are applied for high and low redshift galaxies.  In section \\ref{sec:back-reaction}, the back-reaction to a host galaxy is discussed with respect to the galactic structure and the production of high velocity stars.  In section \\ref{sec:summary}, we summarise this paper. ",
        "conclusions": "\\label{sec:summary}  We have performed $N$-body simulations to investigate successive mergers of MBHs in galaxies with different masses and radii. We have found that about a half of multiple MBHs successively merge to one bigger MBH within $140t_{\\rm dy,g}$ in galaxies with the velocity dispersion larger than $\\sim 180$~km~s$^{-1}$. The merger of MBHs is promoted, such that the loss cone of binary MBHs is refilled by MBHs losing their angular momenta due to dynamical friction. GW recoil does not affect the merger process, if the recoil velocity is of the order of the stellar velocity dispersion.  Galaxies which allow multiple MBHs to merge should reside in dark matter haloes with the mass more than $4 \\times 10^{10}M_{\\odot}$, if these dark matter haloes form at high redshifts. These galaxies could correspond to LAEs or LBGs brighter than the UV magnitude $M_{\\rm UV} \\simeq -19$ at high redshifts.  On the other hand, an MBH which has experienced the successive merger can inhabit low-redshift galaxies brighter than $M_{\\rm UV} \\simeq -18$.  We have also investigated the evolution of the galactic structure and the generation of high-velocity stars as the back-reaction by the successive merger of MBHs.  We have found that the dynamics of MBHs affects the central regions of galaxy that contain about ten times the total mass of MBHs. The mass density profile is transformed to $\\rho \\propto r^{-0.5}$, which is the same as the mass density profile in the case of a galaxy with two and three MBHs. The mass density in the central regions is $1.5$ times smaller than in the case of the galaxy with two MBHs.  In a galaxy with ten MBHs, high-velocity stars are generated intermittently, while they are generated at a constant rate in the case of a galaxy with two MBHs. Such features should enable us to constrain the merger mechanism of MBHs in a galaxy."
    },
    "1402/1402.3614_arXiv.txt": {
        "abstract": "Recently the number of main-sequence and subgiant stars exhibiting solar-like oscillations that are resolved into individual mode frequencies has increased dramatically. While only a few such data sets were available for detailed modeling just a decade ago, the {\\it Kepler} mission has produced suitable observations for hundreds of new targets. This rapid expansion in observational capacity has been accompanied by a shift in analysis and modeling strategies to yield uniform sets of derived stellar properties more quickly and easily. We use previously published asteroseismic and spectroscopic data sets to provide a uniform analysis of 42 solar-type {\\it Kepler} targets from the Asteroseismic Modeling Portal (AMP). We find that fitting the individual frequencies typically doubles the precision of the asteroseismic radius, mass and age compared to grid-based modeling of the global oscillation properties, and improves the precision of the radius and mass by about a factor of three over empirical scaling relations. We demonstrate the utility of the derived properties with several applications. ",
        "introduction": "\\label{SEC1}  It is difficult to overstate the impact of the {\\it Kepler} mission on the observation and analysis of solar-like oscillations in main-sequence and subgiant stars. In a review from just a decade ago, \\cite{bk03} highlighted the tentative detections of individual oscillation frequencies in just a few such stars from ground-based observations, and {\\it Kepler} was not even mentioned. Despite funding issues that delayed the mission from an original deployment date in 2006, {\\it Kepler} finally launched in March 2009 and operated almost flawlessly for more than 4 years, slightly exceeding its design lifetime \\citep{bor10}. The archive of public data now includes nearly uninterrupted observations for many thousands of solar-type stars, including short-cadence data \\citep[58.85~s sampling,][]{gil10} for hundreds of these targets. In the span of a decade, the study of solar-like oscillations has been transformed dramatically \\citep{cm13}.  During the first 10 months of science operations, {\\it Kepler} performed a survey for solar-like oscillations in more than 2000 main-sequence and subgiant stars, yielding detections in more than 500 targets from the 1-month data sets. The initial analysis of this ensemble, using empirical scaling relations \\citep{kb95} to generate estimates of radius and mass, suggested a significant departure from the mass distribution expected from Galactic population synthesis models \\citep{cha11}. Subsequent analysis of the sample, using updated effective temperatures \\citep{pin12} and a substantial grid-based modeling effort, led to more precise estimates of the radii and masses as well as information about the stellar ages \\citep{cha14}. Some of the brightest stars from the survey were subjected to a more detailed analysis, including spectroscopic follow-up to determine more precise atmospheric properties \\citep{bru12} plus the identification and detailed modeling of dozens of oscillation frequencies in each star \\citep{mat12}. These studies gave us a preview of what to expect from the subsequent phase of the mission.  Starting with Quarter 5 (Q5), {\\it Kepler's} short-cadence study of solar-like oscillations transitioned to a specific target phase, where extended observations began for a fixed sample of stars identified during the survey. The target list during this phase gave priority to stars showing oscillations with the highest signal-to-noise ratio (S/N), but it also retained the brightest main-sequence stars cooler than the Sun, where the lower intrinsic oscillation amplitudes yielded relatively weak detections from the survey. From Q5 through the end of the mission (Q17), about 200 of the 512 available short-cadence targets were typically specified by the {\\it Kepler} Asteroseismic Science Consortium \\citep[KASC,][]{kje10} and about half of those were intended for the study of solar-like oscillations.  Just like the exoplanet side of the mission, the KASC team gradually improved the data reduction and analysis methods while additional data swelled the archive. Never in the history of the field had such extended monitoring been possible at all, let alone for such a large sample of stars. As a consequence, the availability of reliable sets of input data for stellar modeling lagged well behind the continually expanding time-series for each star in the archive. This delay was primarily due to the challenge of coordinating the efforts of multiple teams, first to produce optimized light curves from the raw {\\it Kepler} data \\citep{gar11}, then to fit the global oscillation properties and remove the stellar granulation background from the power spectra \\citep{ver11,mat11}, and finally to extract and identify the individual oscillation frequencies using so-called ``peak-bagging'' techniques \\citep{app12}. Also like the exoplanet program, ground-based follow up observations were difficult to obtain for the faintest targets, further limiting the number of stars for which detailed modeling was feasible.  Even after reliable sets of observational constraints became available, an analogous effort was required to consolidate the results from many stellar modeling teams. Initially this effort sought to define objective metrics of model quality, and to use the ensemble of results from different codes and methods to estimate the systematic uncertainties for a few specific targets \\citep{met10,met12,cre12,dog13,sil13}. The first large sample to emerge from the survey made this ``boutique'' modeling approach impractical, and motivated the initial large-scale application of the Asteroseismic Modeling Portal \\citep[AMP,][]{met09,woi09}. \\cite{mat12} presented a uniform analysis of 22 {\\it Kepler} targets observed for 1 month each during the survey phase, and compared detailed modeling from AMP with empirical scaling relations and with results from several grid-based modeling methods that matched the global oscillation properties ($\\Delta\\nu$ and $\\nu_{\\rm max}$, see below) rather than the individual frequencies from peak-bagging. The results clearly demonstrated the improved level of precision that was possible from detailed modeling of the individual oscillation frequencies, particularly for stellar ages.  In this paper we present stellar modeling results from AMP for the first large sample of {\\it Kepler} targets with extended observations during the specific target phase of the mission. In section~\\ref{SEC2} we describe the sample, which was drawn from the most recently published observations. We outline our stellar modeling approach in section~\\ref{SEC3}, including several improvements to the previous version of AMP and using slightly customized procedures for different types of stars. We present the main results and initial applications in section~\\ref{SEC4}, and we discuss conclusions and future prospects in section~\\ref{SEC5}. ",
        "conclusions": "\\label{SEC5}  Our uniform asteroseismic analysis for a large sample of {\\it Kepler} main-sequence and subgiant stars yielded precise determinations of many stellar properties (see Table~\\ref{tab1}), and some important lessons for future work. Most of the simple stars have values of $f\\le10$, indicating a reasonable match between the models from AMP and the observational constraints. The exceptions (KIC\\,6603624, 8694723, 8760414 and 11244118) are stars with the most extreme metallicities in our sample, which may require modifications to the solar-composition mixture in the opacity tables employed by AMP for more accurate modeling. More of the F-like stars have $f>10$, but these are generally targets where the $\\ell=0$ and $\\ell=2$ modes are most difficult to separate. The {\\'e}chelle diagrams in Figure~\\ref{fig3} show that the extracted frequencies for F-like stars from \\cite{app12} include variations in the small separations $d_{0,2}(n)$ that cannot be reproduced by models---in which the $\\ell=0,2$ ridges always curve together with only monotonic variations in $d_{0,2}(n)$. This suggests that the smoothness criteria in peak-bagging pipelines may need to be modified for F-like stars. Automated fitting of the mixed-mode stars is limited to those with fewer than $\\sim$3 avoided crossings and the match to the observed frequencies is generally less precise, with most stars showing $f>10$. The greatest difficulties are again encountered for extreme metallicity targets (KIC\\,7976303 and 10018963), but the qualitative agreement with the locations of avoided crossings for most of the mixed-mode stars in Figure~\\ref{fig3} is remarkable. Ultra-precise constraints on the properties of these stars must rely on future dense grid-modeling to refine the estimates from Table~\\ref{tab1}.  Our modifications to the fitting procedures to try and avoid a bias towards low-helium solutions have helped, but the problem was not eliminated entirely. As noted in section~\\ref{SEC3.1}, of the 22 stars considered by \\cite{mat12} six of them (27\\%) yielded an initial helium mass fraction significantly below the primordial value, while four additional targets (18\\%) were marginally below $Y_{\\rm P}$. We attempted to address this issue by including the frequency ratios as additional constraints, and by adopting larger uncertainties at higher frequencies where the surface correction is larger. A smaller fraction of our sample of 42 stars appears to be affected by the low-helium bias, but Table~\\ref{tab1} still includes one star (2.4\\%) with $Y_{\\rm i}$ significantly below $Y_{\\rm P}$ and three additional targets (7.1\\%) that fall marginally below $Y_{\\rm P}$. Further improvement may require an alternative surface correction \\citep[e.g.,][]{jcd12}.  The precision of asteroseismic data sets from extended observations with {\\it Kepler} now demands that we address the dominant sources of systematic error in the stellar models. The largest source arises from incomplete modeling of the near-surface layers. Although something like the empirical correction of \\cite{kje08} will continue to be useful, we are now in a position to capture more of the relevant physics. For example, \\cite{gru13} performed a Bayesian analysis of the 22 {\\it Kepler} targets from \\cite{mat12}, including a simplified non-adiabatic treatment of the pulsations. They found that the Bayesian probabilities were higher when non-adiabatic rather than adiabatic frequencies were fit to the observations, and that for most stars the surface effect was minimized and in some cases even eliminated. Their non-adiabatic model accounted for radiative losses and gains but neglected perturbations to the convective flux and turbulent pressure \\citep{gue94}. In the case of the Sun, the stability and frequency of the oscillation modes depends substantially on turbulent pressure and the inclusion of non-local effects in the treatment of convection \\citep{bal92a,bal92b,hou10}, and the resulting frequency shift is uncertain. The sensitivity of the results to the model of convection and the temperature profile in the super-adiabatic layer was also emphasized by \\cite{ros99} and \\cite{li02}. Regardless, the Bayesian approach has the advantage of incorporating the unknown sources of systematic error directly into the uncertainties on the derived stellar properties, and can reveal which approach to the pulsation calculations generally improves the model fits.  For future analyses, we intend to augment ADIPLS with the non-adiabatic stellar oscillation code GYRE \\citep{tt13}, which currently includes a limited treatment of non-adiabatic effects but is flexible enough to incorporate additional contributions. We would also like to take advantage of the numerical stability and modular architecture of the open-source MESA code \\citep{pax13} to explore different chemical mixtures and to include heavy element diffusion and settling in the evolutionary models, which is not currently stable for all types of stars with ASTEC. With these new modules for stellar evolution and pulsation calculations, we can embed a Bayesian formalism into the parallel genetic algorithm to complement the simple $\\chi^2$ approach. The complete sample of asteroseismic targets from extended {\\it Kepler} observations spanning up to four years will provide a rich data set to validate these new ingredients for the next generation of AMP.  The golden age of asteroseismology for main-sequence and subgiant stars owes a great debt to the {\\it Kepler} mission, but it promises to continue with the anticipated launch of NASA's Transiting Exoplanet Survey Satellite \\citep[TESS,][]{ric14} in 2017. While {\\it Kepler} was able to provide asteroseismic data for hundreds of targets and could simultaneously monitor 512 stars with 1-minute sampling, TESS plans to observe $\\sim$500,000 of the brightest G- and K-type stars in the sky at a cadence sufficient to detect solar-like oscillations. The data sets will be nearly continuous for at least 27 days, but in two regions near the ecliptic poles the fields will overlap for durations up to a full year. These brighter stars will generally be much better characterized than the {\\it Kepler} targets---with parallaxes from {\\it Hipparcos} and ultimately {\\it Gaia} \\citep{per01}, and reliable atmospheric constraints from ground-based spectroscopy---making asteroseismic characterization more precise and accurate. With several years of development time available, AMP promises to be ready to convert this avalanche of data into reliable inferences on the properties of our solar system's nearest neighbors."
    },
    "1402/1402.4811_arXiv.txt": {
        "abstract": "We present high-energy (3--30~keV) {\\it NuSTAR} observations of the nearest quasar, the ultraluminous infrared galaxy (ULIRG) Markarian~231 (Mrk 231), supplemented with new and simultaneous low-energy (0.5--8~keV) data from {\\it Chandra}.  The source was detected, though at much fainter levels than previously reported, likely due to contamination in the large apertures of previous non-focusing hard X-ray telescopes. The full band (0.5--30 keV) X-ray spectrum suggests the active galactic nucleus (AGN) in Mrk~231 is absorbed by a patchy and Compton-thin (N$_{\\rm H} \\sim1.2^{+0.3}_{-0.3}\\times10^{23}$~cm$^{-2}$) column.  The intrinsic X-ray luminosity (L$_{\\rm 0.5-30~keV}\\sim1.0\\times10^{43}$~erg~s$^{-1}$) is extremely weak relative to the bolometric luminosity where the 2--10~keV to bolometric luminosity ratio is $\\sim$0.03\\% compared to the typical values of  2--15\\%.  Additionally, Mrk~231 has a low X-ray-to-optical power law slope ($\\alpha_{\\rm OX}\\sim-1.7$).  It is a local example of a low-ionization broad absorption line (LoBAL) quasar that is intrinsically X-ray weak.  The weak ionizing continuum may explain the lack of mid-infrared [O IV], [Ne V], and [Ne VI] fine-structure emission lines which are present in sources with otherwise similar AGN properties. We argue that the intrinsic X-ray weakness may be a result of the super-Eddington accretion occurring in the nucleus of this ULIRG, and may also be naturally related to the powerful wind event seen in Mrk 231, a merger remnant escaping from its dusty cocoon. ",
        "introduction": "\\label{sec:intro}    The presence of outflows in luminous quasars and active galactic nuclei (AGN) is most evident in broad absorption line (BAL) systems \\citep{lynds67}.  Approximately 20\\% of quasars show BAL features at some level \\citep[e.g.,][]{gibson09}, understood to be the result of an observational sightline through AGN-driven wind material \\citep[e.g.,][]{weymann91, ogle99, dipompeo13}.  In a common model for BAL quasars \\citep[e.g.,][]{murray95, proga00}, the wind is launched from the accretion disk approximately 0.01--0.1 pc from the black hole and radiatively driven by ultraviolet (UV) line pressure.  However, X-ray emission from typical quasars would over-ionize the gas, quenching the line-driving mechanism and thus preventing the appearance of BAL features.  Therefore, this so-called ``disk-wind'' model requires that BAL quasars have lower levels of X-ray emission, which could either be intrinsic or be due to shielding very close to the black hole, near the base of the wind.  Indeed, BAL systems are very often observed to be X-ray faint \\citep[e.g.,][]{g06, gibson09, wu10}, though observations have yet to provide a clear picture of the origin of this X-ray faintness.  Absorption of a nominal X-ray continuum has been established in several cases \\citep[e.g.,][]{g02, grupe03, shemmer05, giustini08}, but there may be diversity among the population.  Following the terminology of other studies of X-ray emission in BAL quasars \\citep[e.g.,][and references therein]{luo13}, the term ``X-ray weakness or faintness'' refers to the observed X-ray emission being significantly lower than expected given the optical-UV emission.  This may be caused by a normal X-ray continuum modified by absorption from gas and dust.  The term ``intrinsic X-ray weakness'' refers to an innate property of the AGN, which is one possible cause of the observed X-ray faintness.  In order to investigate whether intrinsic X-ray weakness or shielding gas is the most likely explanation for the observed X-ray faintness of BAL quasars, sensitive observations above 10~keV are highly desirable as they can distinguish between the presence of Compton-thick ($N_{\\rm H} \\geq 1.5 \\times 10^{24}$~cm$^{-2}$) shielding material and a less-obscured, intrinsically X-ray weak nucleus.  The {\\it Nuclear Spectroscopic Telescope Array} ({\\it NuSTAR}; \\citealp{nustar}) is the first focusing high-energy X-ray telescope in orbit, and provides unprecedented sensitivity and angular resolution for high-energy, or hard (3--79 keV) X-ray photons.  Recently, {\\it NuSTAR} observed two of the optically brightest BAL quasars known, PG~1004+130 and PG~1700+518, with the goal of investigating the origin of their X-ray faintness \\citep{luo13}.  Unfortunately, low count rates prevented a definitive answer, but they conclusively showed that any absorption present must be Compton-thick.  \\citet{luo13} were also able to demonstrate that about 17--40\\% of BAL quasars are intrinsically X-ray weak by stacking {\\it Chandra} observations of $z\\approx1.5-3$ BAL quasars in the Large Bright Quasar Survey.  Confirmation of intrinsically X-ray weak AGN in some BAL quasars would have important consequences regarding the origin of BAL features.  Instead of invoking shielding gas with an arbitrarily thick column density, intrinsic X-ray weakness would favor a scenario in which the BAL features are coupled to abnormally faint coronal X-ray emission, perhaps due to some mechanism that (periodically) quenches the X-ray corona (see \\S~4.2 of \\citealp{luo13} for discussion).  To this end, we have obtained an extremely sensitive, broadband (0.5--30~keV) X-ray spectrum of Markarian 231 (Mrk 231), the nearest BAL quasar, using {\\it Chandra} and {\\it NuSTAR}.  This paper is organized as follows:  \\S2 details the multi-wavelength properties of Mrk~231; \\S3 presents the new X-ray observations obtained for this study; \\S4 discusses the previously reported hard X-ray spectra of Mrk~231, which are inconsistent with our new data; \\S5 presents our modeling of the new broadband X-ray spectrum, \\S6 discusses our favored scenario, that Mrk~231 is indeed intrinsically X-ray weak; and \\S7 summarizes our results.  Throughout this paper, we adopt $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.27$, and $\\Lambda = 0.73$ \\citep{cosmo}.  Luminosities taken from the literature have been recalculated for our assumed cosmology.  \\begin{figure*}[ht] \\centering \\includegraphics[width=5.5in, angle=270]{mrk231_orientation_pretty.ps} \\caption{Schematic showing the complexity of Mrk~231's structure and our assumed viewing geometry based on multiwavelength data.  The host galaxy is face-on on the plane of the sky with optical emission extending out to a radius of 8~kpc ($9\\farcs8$) \\citep[e.g.,][]{rv11}.  The thin (23~pc) molecular disk is tilted by 10$^\\circ$ with the denser inner region extending from $\\sim$75~pc ($0\\farcs1$) to 1.1~kpc ($1\\farcs3$) and the outer molecular disk extending out to 2.5~kpc ($3\\farcs0$) \\citep{downes98}.  The accretion disk is assumed to be orthogonal to the pc-scale jet, which is assumed to be inclined by 65$^{\\circ}$ (see \\S\\ref{sec:torus}).  In this figure, north is up.} \\label{fig:geometry} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  We have obtained the most sensitive broadband (0.5--30~keV) spectrum of Mrk~231, the closest BAL quasar and an ULIRG,  to date.  {\\it NuSTAR} detects Mrk~231 at a fainter level above 10~keV than previous observations with poorer angular resolution by {\\it BeppoSAX} PDS and {\\it Suzaku} PIN.  Two epochs of {\\it NuSTAR} data separated by seven months show no evidence of hard X-ray variability so contamination is the likely cause of the discrepancy between {\\it NuSTAR} and previous, non-focusing hard X-ray observations.  Analysis of contemporaneous {\\it Chandra} and {\\it NuSTAR} broadband spectra find significant X-ray emission from a powerful circumnuclear starburst.  The direct AGN emission is absorbed and scattered by a patchy torus that is Compton-thin with N$_{\\rm H}\\sim1.2^{+0.3}_{-0.3}\\times10^{23}$~cm$^{-2}$.  The Compton-thin absorption is compatible with the apparent weak Fe~K emission (EW$\\sim$0.2~keV).%  Mrk~231 appears intrinsically X-ray weak.  The absorption-corrected 2--10~keV luminosity is $\\lesssim0.1$\\% of the AGN bolometric luminosity with $\\alpha_{\\rm OX}\\sim-1.7$.  A reason for the low X-ray luminosity could be that the AGN is in a rare ultra soft state similar to those seen in X-ray binaries, though this is inconsistent with the measured $\\Gamma_{\\rm AGN}$.  It is unclear why the AGN is intrinsically X-ray weak, but there are not very many known examples of such sources as points of comparison.  The weak ionizing continuum explains the lack of mid-infrared [O IV], [Ne V], and [Ne VI] fine structure emission lines which are present in sources with otherwise similar AGN properties.  % The intrinsic X-ray weakness may be a result of the super-Eddington accretion occurring in the nucleus of the ULIRG.  These results on Mrk~231 are consistent with its being a merger remnant emerging from its dusty cocoon through a powerful wind event."
    },
    "1402/1402.4350_arXiv.txt": {
        "abstract": "Relativistic contributions to the dynamics of structure formation come in a variety of forms, and can potentially give corrections to the standard picture on typical scales of 100\\,Mpc. These corrections cannot be obtained by Newtonian numerical simulations, so it is important to accurately estimate the magnitude of these relativistic effects.  Density fluctuations couple to produce a background of gravitational waves, which is larger than any primordial background. A similar interaction produces a much larger spectrum of vector modes which represent the frame-dragging rotation of spacetime.  These can change the metric at the percent level in the concordance model at scales below the equality scale. Vector modes modify the lensing of background galaxies by  large-scale structure. This gives in principle the exciting possibility of measuring relativistic frame dragging effects on cosmological scales. The effects of the non-linear tensor and vector modes on the cosmic convergence are computed and compared to first-order lensing contributions from density fluctuations, Doppler lensing, and smaller Sachs-Wolfe effects. The lensing from gravitational waves is negligible so we concentrate on the vector modes. We show the relative importance of this for future surveys such as Euclid and SKA. We find that these non-linear effects only marginally affect the overall weak lensing signal so they can safely be neglected in most analyses, though are still much larger than the linear Sachs-Wolfe terms. The second-order vector contribution can dominate the first-order Doppler lensing term at moderate redshifts and are actually more important for survey geometries like the SKA. ",
        "introduction": "\\label{sec1}  Relativistic corrections to the standard model of cosmology come in a variety of forms, from the altering the dynamics of structure formation to the various effects associated to the interpretation of observations, in particular modifying the propagation of light.  There has been considerable debate as to the importance and amplitude of these effects on the dynamics of the expansion of the universe and the growth of large scale structure (see, e.g., Ref.~\\cite{Clarkson:2011zq} for an overview), and the amplitude and importance of these dynamical effects are still actively debated~\\cite{buchertetal,nobackreac,backsig}. Though subdominant for linear structure formation, relativistic corrections are a generic prediction of General Relativity and are inevitable at a non-linear level through mode-mode coupling. The scalar gravitational potential induces rotational frame-dragging modes in spacetime (so-called vector modes) as well as gravitational waves (tensor modes). Neither of these have counterparts in Newtonian gravity as they both induce a non-zero magnetic Weyl curvature which is absent in Newtonian gravity and difficult to take into account in N-body numerical simulations~\\cite{Bruni:2013mua,Adamek:2013wja}. They therefore serve as an important tool in understanding purely relativistic aspects of structure formation and its observational consequences, as they set a lower limit on the amplitude of relativistic corrections.  On top of dynamical corrections, relativistic effects also induce corrections to the propagation of light since it probes the complete spacetime geometry. This can alter the interpretation of cosmological observations at a level that cannot be neglected in an era of ``precision cosmology''. Provided one works within perturbation theory, the amplitude of these effects is computable and completely fixed once the normalisation of the scalar power spectrum, at the linear level, is determined. For instance, some relativistic effects have been taken into account on the cosmic microwave background~\\cite{pub} and shown to be below the constraints on non-Gaussianity derived by \\textit{Planck}~\\cite{planckXXIX}, but nevertheless in principle detectable on small angular scales, in particular through spectral distortions~\\cite{spectral}. This article focuses on the effect of relativistic corrections on weak lensing observations, focusing mainly on the induced vector mode background. Weak gravitational lensing by the large-scale structure of the Universe has now become a major tool of cosmology~\\cite{revuesWL}, used to study questions ranging from the distribution of dark matter to tests of general relativity~\\cite{testGR}.  The propagation of light in an inhomogeneous universe gives rise to both distortion and magnification induced by gravitational lensing. The effect of non-linear corrections on the Hubble diagram have been considered~\\cite{BenDayan:2012wi,cemuu,Umeh:2012pn,Umeh:2014ana} and shown to be non-negligible given the accuracy of contemporary observations~\\cite{Fleury:2013sna,Fleury:2013uqa,voidslensing,BenDayan:2012ct,BenDayan:2013gc}. In this article we consider the effect on the weak lensing convergence of non-linear effects that induce the existence of a vector and tensor modes background. We compare this to the various contributions to the convergence at first-order~-- the usual integral of the density contrast along the line of sight~\\cite{revuesWL}, the contribution from the Doppler effect which is dominant at low redshifts and large scales~\\cite{Bonvin:2008ni,voidslensing,Bacon:2014uja}, the Integrated Sachs-Wolfe (ISW) and Sachs-Wolfe (SW) terms which are relatively small and mainly neglected when computing cosmic convergence.   The induced background of gravitational waves from scalar-scalar coupling was presented in Ref.~\\cite{Ananda:2006af} during the radiation era, and its present-day spectrum calculated in Ref.~\\cite{Baumann:2007zm}, with shear lensing effects studied in Ref.~\\cite{Sarkar:2008ii}, all following the pioneering analysis of Ref.~\\cite{Mollerach:2003nq}. Surprisingly it was found that the induced gravitational wave background is significantly larger than any primordial background (even for a tensor-scalar ratio $r\\sim0.1$) on intermediate scales of $\\sim$100\\,Mpc, which is around the equality scale, though of course it is much smaller on small scales. Similarly, the induced vector mode background was presented in Refs.~\\cite{Lu:2007cj,Lu:2008ju}, and again a spectrum was found that peaks on 100\\,Mpc scales. Remarkably, however, it was found that the amplitude of the background of vector modes for the metric potential behaves on small scales with the same scaling as the gravitational potential, with nearly 1\\% of its amplitude. While both of these induced degrees of freedom have little effect on the dynamics of structure formation (they cannot directly source the density fluctuation as it is a scalar degree of freedom) they can influence the gravitational lensing produced by large-scale structure. Is it significant, and could it be a new way to detect relativistic aspects of structure formation?  The effects of these contributions on weak lensing convergence predictions are computed in order to understand if they can either be detected or, in the worst case, bias the analysis of future weak lensing experiments, such as Euclid or SKA; {\\em i.e.}, if the interpretation of the observation by assuming that the observed convergence corresponds to the convergence sourced by scalar modes only is an accurate enough assumption or whether some of these effects have to be included in the analysis. This article addresses this question and computes the effect of these two non-linear effects on weak lensing observations by considering second order vector and tensor background. We restrict our analysis to the hypothesis that the Born approximation still holds. In principle, one needs also to take into account second order effects on the geodesic deviation equation~\\cite{Seitz:1994xf, Cooray:2002mj,Dodelson:2005zj,Schaefer:2005up,Shapiro:2006em}, as fully described in Refs.~\\cite{Bernardeau:2009bm,Bernardeau:2011tc}. This effect is however small.  In Section~\\ref{sec2} we describe the vector and gravity waves background induced by the non-linear dynamics and then, in Section~\\ref{sec3}, the computation of the weak lensing power spectra, splitting the effects of the scalar, vector, tensor, Doppler, ISW and SW contributions in order to compare their magnitude. Since the contribution of the tensor modes remains negligible and both ISW and SW being relatively small, we focus in Section~\\ref{sec4} on the vector and Doppler contribution, estimating their magnitude in surveys such as Euclid and SKA. Technical details are gathered in Appendices~\\ref{app0}-\\ref{appC}. ",
        "conclusions": "\\label{sec5}  This article has evaluated the amplitude of relativistic contributions to the weak lensing power spectra. We have considered the gravitational wave and vector mode backgrounds which are sourced at second-order by density perturbations. The amplitude of these backgrounds are completely fixed once the normalisation of the scalar power spectrum in the linear regime is determined. As these are purely relativistic degrees of freedom they set the lower limit for all relativistic effects on cosmological modelling. While the gravitational wave background is very small in relation to the scalars, the vectors, which represent frame dragging in the metric, give corrections to the metric at nearly the percent level. The effect of these contributions on weak lensing convergence predictions have been computed in order to understand if they can either be detected, or bias the analysis of future weak lensing experiments, such as Euclid or SKA. We have compared them to the usual gravitational lensing contribution, the two Sachs-Wolfe contributions as well as the Doppler lensing contribution~\\cite{Bacon:2014uja}.  First, we have shown that even though the non-linear tensor mode background dominates over any possible primordial gravitational wave contribution, its effect on weak lensing is completely negligible, by 10 to 12 order of magnitudes (see Figs.~\\ref{p12} and~\\ref{points}).  Then, we have shown that the vector contribution to the convergence, while small, can dominate over the Doppler lensing at high redshift~-- but there it is swamped by gravitational lensing by density perturbations. We have shown this both for point sources and for two survey geometries. The vectors are actually more important than the Doppler term for SKA-like source distributions on small scales, but not for a Euclid like survey. For both of these surveys the vectors only reach about $10^{-3}$\\% that of the normal gravitational lensing contribution, and so can be safely neglected. Nevertheless, it is interesting that the vector contribution can be as important as some linear terms.  We have also recovered that although the frame dragging effect is small, it becomes more important than both ISW and SW above $\\ell \\ge 50$. This comes to corroborate the fact that for observations, neglecting the 2 first order Sachs-Wolfe terms is a good approximation.  In this analysis, the non-linear effects of the metric perturbations have been described at second order while weak lensing was described assuming that the Born approximation still holds. In principle, one needs also to take into account second order effects on the geodesic deviation equation~\\cite{Seitz:1994xf, Cooray:2002mj,Dodelson:2005zj,Schaefer:2005up,Shapiro:2006em}, as fully described in Refs.~\\cite{Bernardeau:2009bm,Bernardeau:2011tc}.  There are a huge variety of second-order effects which come into the convergence. We have only considered two contributions which arise from non-linear dynamical effects which happen as structure forms. Many contributions appear when calculating the lensing convergence itself~\\cite{Umeh:2012pn,Umeh:2014ana,BenDayan:2012wi}, and these also need to be analysed in a similar manner to that presented here to determine whether relativistic effects are important for future observations of magnification."
    },
    "1402/1402.4485_arXiv.txt": {
        "abstract": "{We present the most extensive analysis of Fourier-based \\mbox{X-ray} timing properties of the black hole binary Cygnus X-1 to date, based on 12 years of bi-weekly monitoring with RXTE from 1999 to 2011. Our aim is a comprehensive study of timing behavior across all spectral states, including the elusive transitions and extreme hard and soft states. We discuss the dependence of the timing properties on spectral shape and photon energy, and study correlations between Fourier-frequency dependent coherence and time lags with features in the power spectra. Our main results are: (a) The fractional rms in the 0.125--256\\,Hz range in different spectral states shows complex behavior that depends on the energy range considered. It reaches its maximum not in the hard state, but in the soft state in the Comptonized tail above 10\\,keV. (b) The shape of power spectra in hard and intermediate states and the normalization in the soft state are strongly energy dependent in the 2.1--15\\,keV range. This emphasizes the need for an energy-dependent treatment of power spectra and a careful consideration of energy- and mass-scaling when comparing the variability of different source types, e.g., black hole binaries and AGN. PSDs during extremely hard and extremely soft states can be easily confused for energies above $\\sim$5\\,keV in the 0.125--256\\,Hz range. (c) The coherence between energy bands drops during transitions from the intermediate into the soft state but recovers in the soft state. (d) The time lag spectra in soft and intermediate states show distinct features at frequencies related to the frequencies of the main variability components seen in the power spectra and show the same shift to higher frequencies as the source softens. Our results constitute a template for other sources and for physical models for the origin of the \\mbox{X-ray} variability. In particular, we discuss how the timing properties of \\mbox{Cyg~X-1} can be used to assess the evolution of variability with spectral shape in other black hole binaries. Our results suggest that none of the available theoretical models can explain the full complexity of \\mbox{X-ray} timing behavior of \\mbox{Cyg~X-1}, although several ansatzes with different physical assumptions are promising.} ",
        "introduction": "The canonical states of accreting black hole binaries have first been observed and defined in the spectral domain: a hard state with a power law spectrum with a photon spectral index of $\\sim$1.7 and a soft state with a spectrum dominated by thermal emission from an accretion disk. They are joined by usually short-lived transitional or intermediate states.  The whole sequence of states -- from hard state over the intermediate into soft and then again into intermediate and finally into hard state -- can best be depicted on a hardness-intensity-diagram (HID), where transient sources that undergo a full outburst follow a \\textsf{q}-shaped track \\citep[][see \\citeauthor*{McClintock_Remillard_2006a}, \\citeyear{McClintock_Remillard_2006a}, for a different nomenclature]{Fender_2004a}. Radio emission is detected in the hard state, with jets imaged for \\object{\\mbox{Cyg~X-1}} \\citep{Stirling_2001a} and \\object{GRS~1915$+$105} \\citep{Dhawan_2000a,Fuchs_2003a}. In the soft state, radio emission is strongly quenched. Evidence for similar spectral behavior also exists in several other classes of accreting objects such as neutron star \\mbox{X-ray} binaries \\citep[e.g.,][\\object{Aql~X-1}]{Maitra_2004a}, active galactic nuclei \\citep[AGN;][]{Koerding_2006a}, and dwarf novae \\citep[][\\object{SS~Cyg}]{Koerding_2008a}.  The spectral states, including the different flavors of the intermediate state, show distinct \\mbox{X-ray} timing characteristics, such as shapes of power spectra or time lags between emission at different energies.  Timing parameters seem to be a remarkably sensitive tool for defining state transitions \\citep[e.g.,][]{Pottschmidt_2003b,Fender_2009a,Belloni_2010a}.  While the radio emission, and possibly also the gamma-ray emission above 400\\,keV \\citep{Laurent_2011a,Jourdain_2012a} originate in the jet, the origin of the \\mbox{X-ray}s is still unclear. As shown, e.g., by \\citet{Nowak_2011a}, the combination of the best resolution and the most broadband \\mbox{X-ray} spectra available today fails to enable us to statistically distinguish between jet models \\citep{Markoff_2005a, Maitra_2009a} and thermal and/or hybrid Comptonization in a corona \\citep[e.g.,][]{Coppi_1999a,Coppi_2004a}. The fluorescent Fe K$\\alpha$ line and a reflection hump point towards a contribution by reflection, independent of the origin of the continuum \\citep[see][for a review and \\citealt{Duro_2011a}, \\citealt{Tomsick_2014a}, and references therein for \\mbox{Cyg~X-1}]{Reynolds_2003a}.  Spectro-timing analysis holds the promise of solving this ambiguity, as a truly physical model has to consistently describe both spectral and timing behavior. While no self-consistent models that would encompass all parameters exist yet, some studies address, e.g., simultaneous modeling of photon spectra and root mean square variability (rms) spectra \\citep{Gierlinski_2005a} or photon spectra and Fourier-dependent time lags \\citep{Cassatella_2012a}. Theoretical ansatzes to describe the \\mbox{X-ray} variability include propagating mass accretion rate fluctuations \\citep[usually based on][see, e.g., \\citealt{Ingram_2013a} for an analytical model]{Lyubarskii_1997a}, up-scattering in a jet \\citep{Reig_2003a,Kylafis_2008a}, and full magnetohydrodynamic simulations \\citep[][who concentrate on spectra but also address timing properties]{Schnittman_2013a}.  There is mounting evidence for similarities in timing behavior between \\mbox{X-ray} binaries, AGN \\citep[see][for a review]{McHardy_2010a}, and recently also cataclysmic variables \\citep[see][for discovery of Fourier-dependent time lags]{Scaringi_2013a}. Because of their brightness and short variability time scales, \\mbox{X-ray} binaries remain the best laboratories to investigate these phenomena and are key to deciphering the complex and currently highly disputed interplay of accretion and ejection processes. We note especially that AGN are usually viewed on a single state due to the very much longer variation timescale, and thus \\mbox{X-ray} binaries are needed to study the variety of states and their inter-relationships.  The first step to understanding the \\mbox{X-ray} timing characteristics is a fundamental overview of their evolution with the spectral shape that can only be achieved with a high number of high quality observations densely covering all states, including the elusive transitions. Most previous works concentrate on energy-independent evolution of rms and power spectral distributions (PSDs) with spectral state \\citep[e.g.,][]{Pottschmidt_2003b,Belloni_2005a,Axelsson_2006a,Klein-Wolt_2008a}, although further spectral shape-, energy-, and Fourier frequency-dependent correlations have been noted in individual observations and smaller samples \\citep[e.g.,][]{Homan_2001a,Rodriguez_2002b,Kalemci_2004a,Rodriguez_2004a, Boeck_2011a,Cassatella_2012b,Stiele_2013a}, often with a focus on the behavior of narrow quasi-periodic oscillations.  Missing are consistent analyses of the energy-resolved evolution of rms and PSDs, Fourier-frequency dependent evolution of cross-spectral quantities (coherence function and lags), and correlations of features in PSDs and Fourier-frequency dependent cross-spectral quantities over the full range of spectral states and over multiple transitions. In this paper, we address these questions with an extraordinarily long and well sampled RXTE campaign on the high mass black hole binary \\mbox{Cyg~X-1} that enables us to conduct the most comprehensive spectro-timing analysis of a black hole binary to date.  Located at a distance of $1.86^{+0.12}_{-0.11}$\\,kpc \\citep[][consistent with \\citealt{Xiang_2011a}]{Reid_2011a}, \\mbox{Cyg~X-1} is bright (in the hard state $\\sim$$7\\times10^{-9}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ in the 1.5--12\\,keV band of RXTE-ASM) and persistent and therefore a prime target for both spectral and timing studies. Often considered a prototypical black hole binary, it is frequently used for comparisons with other black hole binaries \\citep[e.g,][]{Munoz-Darias_2010a} and AGN \\citep{Markowitz_2003a,McHardy_2004a,Papadakis_2009a}.  Although the spectrum of the source is never fully dominated by the disk and the bolometric luminosity only changes by a factor of $\\sim$4 \\citep[][and references therein]{Cui_1997a,Shaposhnikov_2006a,Wilms_2006a}, \\mbox{Cyg~X-1} often undergoes state transitions \\citep[e.g.,][]{Pottschmidt_2003b,Wilms_2006a,Grinberg_2013a} and is therefore also well suited for studies of the intermediate states.  Here, we analyze data from the full set of RXTE observations of \\mbox{Cyg~X-1}. This paper is a part of a series where we previously analyzed spectro-timing correlation in the hard state 1998 to 2001 \\citep{Pottschmidt_2003b}, the rms-flux relation \\citep{Gleissner_2004a}, the radio-\\mbox{X-ray} correlations \\citep{Gleissner_2004b}, the spectral evolution 1999--2004 \\citep{Wilms_2006a}, and states and state transitions 1996--2012 with all sky monitors \\citep{Grinberg_2013a}. We start Sect.~\\ref{sect:data} by introducing the data and the general behavior of the source during the time period covered by this analysis and follow with a description of the spectral analysis and employed \\mbox{X-ray} timing techniques. In Sect.~\\ref{sect:rms_and_power}, we discuss the energy independent and dependent evolution of the rms and the PSDs with spectral shape. In Sect.~\\ref{sect:cross}, we discuss the evolution of cross-spectral quantities with spectral shape. We address the implication of our results for the analysis of other sources in Sect.~\\ref{sect:other} and for theories explaining the origin of \\mbox{X-ray} variability in Sect.~\\ref{sect:theory}. Sect.~\\ref{sect:summary} summarizes our finding with a focus on the implied directions for further investigations. ",
        "conclusions": "\\label{sect:summary}  In the following we briefly summarize the highlight results of this paper that we expect to have the strongest constraints on future theoretical models and to offer the strongest stimuli for future investigations of the variability of \\mbox{X-ray} binaries: \\begin{itemize} \\item The fractional rms shows a complex behavior that strongly depends on the spectral state of the source and on the considered energy range. In particular, the fractional variability in the soft state is low at energies below $\\sim$4.5\\,keV ($\\sim$15\\%) but increases with energy up to $\\sim$40\\% in the 9.4--15\\,keV range. The ratio of fractional rms at different energies shows a tight dependence on spectral shape. \\item The shape of the power spectra is highly dependent on the energy band used, with a striking change of behavior at $\\Gamma_1 \\approx 2.6$--2.7. This reiterates the importance of an energy-dependent approach to PSDs and of carefully addressing the scaling of typical energies with mass when comparing variability from different source types. \\item The power spectra in hard and intermediate states are dominated by two main variability components that shift to higher frequencies as the source spectrum softens. In the $\\Gamma_1 \\approx 2.4$--2.7 range, the 9.4--15\\,keV band is dominated by the higher frequency component, while in the 2.1--4.5\\,keV and the 4.5--5.7\\,keV bands the lower frequency component is more prominent. \\item The power spectra of the hardest states show further variability components at higher Fourier frequencies. In the 0.125--256\\,Hz range, these PSDs can easily be confused with the softest observations. \\item The coherence is close to unity at frequencies below $\\sim$10\\,Hz in the hard and soft states. At the transition from intermediate to soft state, the coherence drops strongly, with hints of structure similar to the two dominant components of the power spectra. It recovers in the soft state. Above $\\sim$10\\,Hz, the coherence cannot be constrained in our orbit-wise approach. \\item The time lag spectra in the hard and intermediate states show features with a double-humped structure similar to the two dominant components of the PSDs and with the same evolution to higher frequencies as the photon spectra soften. No structure is visible in the time lag spectra with soft photon spectra ($\\Gamma_1 \\gtrsim 2.7$). \\item The average time lags in the 3.2--10\\,Hz band show a non-linear increase with $\\Gamma_1$ until $\\Gamma_1 \\approx 2.65$. For $\\Gamma_1 \\gtrsim 2.65$, the average lags are low and show no dependence on spectral shape. \\end{itemize} We have shown that taking into account spectro-timing correlation does not only provide a more holistic description of source properties, but can also help to assess the quality of spectral models. The spectro-timing analysis presented here for the exceptional data set of \\mbox{Cyg~X-1} can be used as a template for other sources with a worse coverage of different spectral shapes. The behavior it describes will allow to test future theoretical models for accretion and ejection processes against observations. Better constraints of some elusive timing parameters, such as the detailed structure of time lag spectra in individual observations, require larger area instruments and further monitoring of the long-term evolution as may be provided with future satellites such as ASTROSAT \\citep{Paul_2013a} and NICER \\citep{Gendreau_2012a}."
    },
    "1402/1402.4166_arXiv.txt": {
        "abstract": "In this paper I present dynamic models of the radio source Centaurus A, and critique possible models of in situ particle reacceleration (ISR) within the radio lobes.  The radio and $\\g$-ray data require neither  homogeneous plasma nor quasi-equipartition between plasma  and magnetic field; inhomogeneous models containing both high-field and low-field regions are equally likely. Cen A cannot be as young as the radiative lifetimes of its relativistic electrons, which range from a few to several tens of Myr.  Two classes of dynamic models -- flow driven and magnetically driven -- are consistent with current observations; each requires Cen A to be on the order of a Gyr old.  Thus, ongoing ISR must be occuring within the radio source.  Alfven-wave ISR is probably occuring throughout the source, and may be  responsible for maintaining the $\\g$-ray-loud electrons.  It is likely to be supplemented by shock or reconnection ISR which maintains the radio-loud electrons in high-field regions. ",
        "introduction": "The full story of the radio galaxy Centaurus A is still not known, despite much careful study over the years.  Cen A is of course worth study for its own sake, but it also provides a nearby, well-documented example from which we can understand the dynamic evolution of other, similar radio galaxies. In this paper I focus on two unresolved questions related to Cen A's large-scale radio lobes.  The first question is the age of Cen A.  To answer this we need to understand the physical state of its radio lobes. So-called Fanaroff-Riley Type II (``FRII'') radio galaxies are well understood: a collimated, supersonic jet develops a terminal shock where it impacts the ambient medium.  The shock is apparent as an outer hot spot, and jet plasma which has been through the shock is left behind as a cocoon.  Ram pressure drives the growth of the source, as the end of the jet advances into the ambient medium. However, Cen A is not an FRII, but a Fanaroff-Riley Type I  (``FRI''; implicitly defined these days as any radio galaxy that is not an FRII).  No robust models  exist yet for the dynamics and development of FRI sources.  The inner jets in some FRI's have been modelled as turbulent, entraining flows, but that model does not address the variety of FRI morphologies, nor does it describe the history and growth of such a source.   More work is needed before we can understand the dynamics of FRI sources.  The second question is how Cen A keeps shining long past the radiative lifetimes of its radio-loud and $\\g$-ray loud electrons. In this paper I discuss three alternative dynamical models which may describe Cen A.  All of these models require the source to be quite old, with an age between several hundred Myr to more than a Gyr. This age range is much longer than the electrons' radiative lifetimes, which range from a few to a few tens of Myr. It follows that the radiative lifetimes are not a good indicator of the age of the source;  in situ reacceleration (ISR) must be ongoing within the large radio lobes.  Thus, Cen A provides a nearby case study of ISR in radio galaxies.     To address these questions, I begin in Section  \\ref{BasicStuff} by describing the setting and collecting important facts to be used in the analysis. In  Section \\ref{InternalConditions} I review observational constraints on plasma and magnetic field within the outer lobes.  I argue that inhomogeneous models, containing both strong and weak fields, are physically likely and consistent with the data.  In Section \\ref{ModelsThatDontWork} I show that the short radiative lifetimes require ISR;  models fail which try to grow the source so quickly. In Section  \\ref{ModelsThatWork} I present three alternative dynamic models -- buoyant, magnetic tower and flow-driven -- each of which requires the source age to be on the order of a Gyr. The buoyant model does not fare well when compared to the data but both the flow-driven and magnetically-driven models seem consistent with current observations.  In Section  \\ref{InSituReacceleration} I critique competing ISR  models. I speculate that Alfven-wave ISR is operating throughout the lobes and probably maintains the $\\g$-ray-loud electrons, while the radio-loud electrons are energized by shock or reconnection ISR in high-field regions within the source. Finally in Section  \\ref{Conclusions} I summarize my arguments and discuss what this exercise reveals about the astrophysics of Cen A. I relegate necessary but tedious details to the Appendices.  \\begin{figure*}[htb] {\\center \\includegraphics[width=0.95\\columnwidth]{EilekFigure1.eps}} \\caption{The large-scale structure of Centaurus A, displayed with north to the right.  Neither lobe shows any sign of limb brightening; the volume of each lobe appears to be filled with localized, radio-bright filaments and loops.  Elliptical rings in the center of this image are not physical, but are part of the interferometer response (to the very strong central source) which has not been fully removed from the image.   The inner lobes  sit within and are obscured by these artifacts.  Mosaiced 1.4-GHz image using ATCA and Parkes data, $60'' \\times 40''$ resolution (data are publicly available on the NASA Extragalactic Database;  see also Feain \\etal 2011).} \\label{Fig:OuterLobes} \\end{figure*} ",
        "conclusions": "\\label{Conclusions}  In this section I summarize my previous arguments in this paper, discuss which of the models I have proposed might best describe Cen A, and speculate on how these models connect to other evidence we have on the recent history of the galaxy.  \\subsection{Summary I:  tentative  models for Cen A}  In this paper I have shown that the radio galaxy Cen A is on the order of 1 Gyr old;  substantially younger models violate observations.  Because this is much older than the radiative lifetimes of the relativistic electrons seen in radio and $\\g$-rays, in situ reacceration (ISR) must be ongoing throughout the 300-kpc-scale radio lobes. I developed the argument in three parts, as follows.  $\\bullet$ I showed (in Section \\ref{InternalConditions}, also Table \\ref{OL_Bfields}) that homogeneous, quasi-equipartion ($u_B \\sim u_e$) models are {\\it not} required  by the data. Inhomogeneous models, in which high-B and low-B regions coexist, are also consistent with the data and probably closer to the real physics.  The high-field regions could be shocks within the outer lobes or force-free regions within a large-scale magnetic structure.   $\\bullet$ I introduced three toy models to describe the dynamical growth of Cen A (in Section \\ref{ModelsThatWork}, also Tables \\ref{Model_summary} and \\ref{Model_critique}). (1) The slowest-growing model assumes the outer ends of the radio lobes are driven  by {\\it buoyancy}.  (2) A second model assumes the radio lobes are magnetically dominated, with large-scale organized structures akin to {\\it magnetic towers}.  (3) A third, faster-growing model assumes radio lobe growth is driven by {\\it internal plasma flow}.  These models require Cen A to be at least several hundred Myr old, and possibly older than a Gyr.  By comparison, the radiative lifetimes range from a few Myr for $\\g$-ray-loud electrons, to at most $\\sim 80$ Myr for radio-loud electrons; thus ISR is needed.   $\\bullet$ I explored (in Section \\ref{InSituReacceleration}, also Table \\ref{Model_critique}) three competing ISR models relevant to Cen A. (1) {\\it Alfven wave turbulence} can explain the data if the radio lobe plasma  is tenuous and highly variable in time and/or space. (2)  {\\it Weak shocks} are an  attractive ISR mechanism in the transonic, flow-driven model; they may account for most or all of the radio emission in that model. (3)  {\\it Reconnection} is another attractive ISR mechanism, especially in the magnetically driven model, {\\it if} reconnection site physics is consistent with  current  thinking.   \\subsection{Summary II: which models might work for Cen A?}  Having introduced three dynamic models, and explored what ISR mechanism(s) can work within each one, I then compared the models to the data (in Section \\ref{Compare_dynamic_models} and Section \\ref{Compare_ISR_models};  also Table \\ref{Model_critique}).  While the simplest form of the buoyant model had severe challenges, two models remain viable: flow driven and magnetically driven.   \\begin{itemize}   \\item  {\\em Flow-driven models} are the closest to traditional models of FRII sources, and may provide a good description of many tailed FRI sources. Their two-part structure (channel flow plus surrounding cocoon) gives a natural explanation for the transverse spectral structure seen in Cen A. If the channel flow is transonic, the internal shocks it drives are an attractive  way to explain the radio filaments and radio spectrum of Cen A.   \\item {\\em Magnetically driven models}, while less popular in the radio galaxy literature,  are an intriguing alternative to flow-driven models. The large-scale, self-organized  structures inherent in these models are a tempting explanation for the large structures and gradients seen in total and polarized radio emission. If these models describe Cen A, large-scale reconnection sites probably supplement Alfvenic ISR throughout the lobes.  \\end{itemize}  However, neither of these models has been developed to the point where it can confidently be said to describe Cen A. Until my suggested scenarios in this paper are supported by dedicated modelling and/or simulations, they remain no more than speculations. While the internal physics of competing ISR mechanisms clearly needs further work, the most immediate -- and addressible -- test of each model  is under what conditions it can explain the large-scale structure of Cen A.  Flow-driven models are perhaps most challenged by the lack of brightening at the outer ends of the lobes, where the channel flow should be decelerating and/or developing terminal shocks.  I have speculated that lower Mach number flows will not suffer such problems, but -- because most work on flow-driven models has considered high Mach number flows -- it is not yet clear under what conditions, if at all,  models driven by slower flows can avoid disruptive instabilities and create large-scale radio lobes resembling those in Cen A.  Magnetically-driven models are perhaps most challenged by the lack of large, synchrotron-faint internal cavities as predicted by existing models.  I have suggested that the radio lobes may contain more complex, self-organized magnetic structures, but this remains to be demonstrated. Work is especially needed to learn whether AGN-driven plasmas can access the necessary high-order magnetic modes in a soft-boundary system (such as a radio lobe in the IGM), and whether the plasma/field mix inside such systems is sufficiently radio-loud to create radio lobes resembling those in Cen A.     \\subsection{Connect  to the recent history of the galaxy} \\label{RecentHistory}  To close, some possible connections between Cen A and recent events in the history of its parent galaxy deserve mention.   \\paragraph{Very recent history:  disrupted  jets?}   Although active jets from the AGN are feeding the kpc-scale inner radio lobes, no detectable jets connect the inner lobes to the outer lobes on either side of the AGN (Neff \\etal 2014). If ongoing energy transport requires a detectable jet, then both outer lobes are disconnected from their engine.  Are they dying?  If so, how quickly will they fade from view?    Because the radio- and $\\g$-ray-loud electrons require frequent ISR in order to keep shining, we must ask, on what timescale the ISR will cease if energy flow to the radio lobes is disrupted.  Turbulence is the common denominator in all three ISR mechanisms.  It is clearly required for Alfven-wave ISR, and it is implicated in our two alternative ISR models.  Thus, the lobes will fade when the turbulence decays.  We know that nondriven fluid turbulence will decay in no more than several turnover times, $\\tau_t = \\lambda_t/v_t$ (where $\\lambda_t$ and $v_t$ are the characteristic scale and velocity of the turbulence. We can estimate $\\lambda_t \\sim \\lambda_{max} \\sim 30$ kpc from the images,  but we have no direct probe of $v_t$.  The flow-driven model allows a possible {\\it ansatz}:  if the field is maintained by a turbulent dynamo,  $B^2 / 4 \\pi \\ltw \\rho v_t^2$, and thus $v_t \\gtw v_A$.  Alternatively, with the magnetically driven model we can directly take $v_A$ as the turbulent turnover speed. Thus, our  limit $v_A \\gtw 4000$ km/s (from Section \\ref{AlfvenISR}) suggests $\\tau_t \\ltw 8$ Myr.  We therefore expect turbulence in the outer lobes to decay in no more than $\\sim 30$ Myr. This is longer than the lifetime of the $\\g$-ray-loud electrons, and  comparable to the lifetime of the radio-loud electrons;  the lobes will fade quickly if the turbulence stops being driven.  It follows that, if the outer lobes are disconnected from the AGN right now, this situation cannot have lasted long. Whatever disrupted the energy flow from the AGN must have happened no more than a few tens of Myr ago. Did the AGN go through a quiescent period, for the past few tens of Myr, and then revive itself only a few Myr ago?  If so, was this part of a regular duty cycle, or is it evidence of an unusual disruptive event in the galaxy's recent history?   \\paragraph{Previous history:  did a merger launch the radio galaxy?}  We have seen that Cen A is  on the order of a Gyr old. It is interesting to compare this age range to the dynamic history of NGC 5128. The evidence suggests that at least two significant events -- encounters or megers -- have  happened to this galaxy in the past few Gyr.   One  line of evidence comes from the dramatic gas/dust disk in the core of NGC5128. Because it is warped and twisted  we must be seeing it fairly soon after its formation, as it settles into a stable state within the galaxy's gravitational potential. The exact rate at which the disk settles depends on details of the galaxy's structure.  If the galaxy is prolate, models suggest the disk is only $\\sim 200$ Myr old (Quillen \\etal 1993);  if the galaxy is oblate, the disk is $\\sim 700$ Myr old (Sparke 1996). If, as generally assumed, the disk has been left behind by a recent merger, that event happened less than a Gyr ago.  Another line of evidence  comes from the galaxy's stars; NGC5128 has a complex star formation history.  While the bulk of the stars are at least $10$ Gyr old, recent work has found a younger population, which may be as young as $\\sim 2\\!-\\!3$ Gyr (\\eg, integrated light spectroscopy of globular clusters, Woodley \\etal 2010; resolved studies of halo stars, Rejkuba \\etal 2011).  If this recent burst of star formation was triggered by an encounter or merger, that event happened $\\sim 2\\!-\\!3$ Gyr ago.   While we do not know what initiates an activity cycle in an AGN, the ages of these two recent events are intriguing. The (relatively young) age of the disk  and the (somewhat older) age of the young stellar population in NGC5128 bracket the ages that the dynamic models predict for the outer lobes. Did one of these recent encounters also trigger AGN activity and create the radio galaxy Cen A?  \\ack  It is a pleasure to thank Susan Neff,  Frazer Owen, Sarka Wykes, Hui Li and Martin Hardcastle for many illuminating conversations about radio galaxy physics in general, and Cen A in particular.  Insightful comments from the referees contributed significantly to this paper. Norbert Junkes kindly provided his Parkes image of Cen A. This research 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. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities Inc.     \\begin{appendix}"
    },
    "1402/1402.6449_arXiv.txt": {
        "abstract": "We construct a self-interacting scalar dark matter (DM) model with local discrete $Z_{3}$ symmetry that stabilizes  a weak scale scalar dark matter $X$. The model assumes a hidden sector with a local $U(1)_X$ dark gauge symmetry, which is broken spontaneously into $Z_3$ subgroup by nonzero VEV of dark Higgs field $\\phi_X$ ($ \\langle \\phi_X \\rangle \\neq 0$).  Compared with global $Z_3$ DM models, the local $Z_3$ model has two new extra fields: a dark gauge field $Z^{'}$ and a dark Higgs field $\\phi$ (a remnant of the $U(1)_X$ breaking).  After imposing various constraints including  the upper bounds on the spin-independent direct detection cross section and thermal relic density,   we find that the scalar DM  with mass less than $125$ GeV is allowed in the local $Z_3$  model, in contrary to the global $Z_3$ model. This is due to new channels in the DM pair annihilations open into $Z^{'}$ and $\\phi$ in the local $Z_3$ model. Most parts of the newly open DM mass region can be probed by XENON1T and other similar future experiments. Also if $\\phi$ is light enough  (a few MeV $\\lesssim m_\\phi \\lesssim$ O(100) MeV), it can generate a right size of DM self-interaction and explain the astrophysical small scale structure anomalies. This would lead to exotic decays of Higgs boson into a pair of dark Higgs bosons, which could be tested at LHC and ILC. ",
        "introduction": "Although Planck \\cite{Ade:2013zuv} has already given the dark matter(DM) relic density $\\Omega h^2=0.1199\\pm 0.0027$ with a high precision, we still do not know particle physics nature of DM at all.   So far all the compelling evidences for the existence of DM come from astrophysics and cosmology, due to its gravitational interaction. Still, many particle physics models for DMs have been proposed, and most of them have a stable collisionless cold DM(CCDM) candidate whose self-interaction can be ignored.  The collisionless cold DM has been very successful when explaining the large scale structure of our Universe. However, anomalies from the small scale astrophysical observations~\\cite{Oh:2010ea,BoylanKolchin:2011de,BoylanKolchin:2011dk} indicate that DM may have strong interactions between themselves. Such self-interaction~\\cite{Spergel:1999mh} would make DM have a flat core density profile rather than a cusp one predicted by CCDM. Recent simulations show that in order to flatten the cores of galaxies the cross section for DM scattering should be around $\\sigma\\sim  M_X \\times \\textrm{ barn }\\GeV^{-1}$~\\cite{Vogelsberger:2012ku, Rocha:2012jg, Zavala:2012us},  which is in fact a huge cross section compared with typical weak-scale cross sections $\\sigma\\sim 10^{-12}$ barn or $1$ pb.  Some light particle mediator in the dark sector could be an origin of such strong self-interaction between DMs.  In this paper, we propose a scalar DM model with a local $Z_3$ symmetry. Unlike models based on global symmetries, local discrete symmetries can protect symmetry-breaking from quantum gravity effects and guarantee the longevity or absolute stability of DM particles. Also a light mediator can exist in the models with local symmetry, and generate the correct self-interaction for DM in explaining the anomalies mentioned in  the previous paragraph.  The outline of this paper is as follows.  In Sec.~II, we introduce the model with a local $Z_3$ symmetry, establish the convention for parameters and give the physical mass spectra.  Then we discuss both theoretical and experimental constraints on the parameters in Sec.~III.  Then in Sec.~IV, we discuss the relic density and DM direct searches, paying attentions to the semi-annihilation feature, and compare with the global $Z_3$ mode.  In Sec.~V, we show that a light scalar mediator in our model can induce strong interaction for DM.  Finally we summarize the results in Sec.~VI. ",
        "conclusions": "In this paper, we have proposed a self-interacting scalar DM model with a local dark $Z_3$ symmetry. Unlike global dark symmetries, local ones can guarantee that DM is absolutely stable even in the presence of higher dimensional nonrenormalizable operators due to the underlying local gauge symmetry.   Then we discussed perturbativity constraints on the scalar potential and the experimental limit on the kinetic mixing. Compared with a global $Z_3$ model, our scenario has two new particles, $Z^{'}$ and $H_2$,  and there are new channels in the DM pair annihilations for thermalizing DMs. Therefore much ampler parameter space is allowed including a light DM with $M_X<125$ GeV,  most region of which can be probed with future DM direct searches. Also, motivated by the small scale astrophysical anomalies, we investigated the phenomenology of a MeV scalar $H_2$ in our model which has no counterpart in the minimal global $Z_3$ model. Thanks to the velocity dependence of DM self-interaction cross section,  such a light $H_2$ can mediate strong interaction for DM scattering at Dwarf galaxy scale while satisfying Milky Way and cluster scale constraints. Similar arguments go for the light $Z^{'}$ as well.  For such a light $H_2$ or $Z^{'}$, there could be exotic decays of the 126 GeV Higgs boson, which could be studied in the upcoming LHC running and at future lepton colliders."
    },
    "1402/1402.0868_arXiv.txt": {
        "abstract": "{} { We aim to investigate cool-core and non-cool-core properties of galaxy groups through X-ray data and compare them to the AGN radio output to understand the network of intracluster medium (ICM) cooling and feedback by supermassive black holes. We also aim to investigate the brightest cluster galaxies (BCGs) to see how they are affected by cooling and heating processes, and compare the properties of groups to those of clusters.}  {Using Chandra data for a sample of 26 galaxy groups, we constrained the central cooling times (CCTs) of the ICM and classified the groups as strong cool-core (SCC), weak cool-core (WCC), and non-cool-core (NCC) based on their CCTs. The total radio luminosity of the BCG was obtained using radio catalogue data and/or literature, which in turn was compared to the cooling time of the ICM to understand the link between gas cooling and radio output. We determined K-band luminosities of the BCG with 2MASS data, and used a scaling relation to constrain the masses of the supermassive black holes, which were then compared to the radio output. We also tested for correlations between the BCG luminosity and the overall X-ray luminosity and mass of the group. The results obtained for the group sample were also compared to previous results for clusters.}  {The observed cool-core/non-cool-core fractions for groups are comparable to those of clusters. However, notable differences are seen: 1)~for clusters, all SCCs have a central temperature drop, but for groups this is not the case as some have centrally rising temperature profiles despite very short cooling times; 2)~while for the cluster sample, all SCC clusters have a central radio source as opposed to only 45\\% of the NCCs, for the group sample, \\textit{all} NCC groups have a central radio source as opposed to 77\\% of the SCC groups; 3)~for clusters, there are indications of an anticorrelation trend between radio luminosity and CCT. However, for groups this trend is absent; 4)~the Indication of a trend of radio luminosity with black hole mass observed in SCC clusters is absent for groups; and 5)~similarly, the strong correlation observed between the BCG luminosity and the cluster X-ray luminosity/cluster mass weakens significantly for groups.}  {We conclude that there are important differences between clusters and groups within the ICM cooling/AGN feedback paradigm and speculate that more gas is fueling star formation in groups, than in clusters where much of the gas is thought to feed the central AGN.}  \\titlerunning{ICM Cooling, AGN Feedback and BCG properties of Galaxy Groups} \\authorrunning{V.~Bharadwaj et al.} ",
        "introduction": "The central entropy is another parameter often used to classify CC/NCC clusters (e.g.~\\citealt{ 2008ApJ...687..899R}) and displays a tight correlation with the CCT (\\citealt{ 2010A&A...513A..37H}). This is not surprising since the cooling time, $t_{\\mathrm{cool}}$, is related to the gas entropy, $K$, through the relation $t_{\\mathrm{cool}} \\propto K^{3/2}/T$ for pure Bremsstrahlung. We also show a histogram of the distribution of the central entropy in Fig.~\\ref{fig:Khist}. The entropy values are given in Table~\\ref{CCT}.  The plot of the CCT and $K_{0}$ values is shown below in Fig.~\\ref{fig:entropy}. As expected, we see an excellent correlation between the two quantities with a Spearman rank correlation coefficient of 0.99, 0.99, and 0.96 for all, SCC, and non-SCC groups respectively. A CCT of 1 Gyr corresponds to $\\approx~20~\\mathrm{keV~cm}^{2}$. \\begin{figure} \\includegraphics[trim = 30mm 0mm 22mm 8mm, clip, scale=0.33]{./Entropyhistgraylogspace_3107.eps} \\caption{Histogram of the central entropy. Light grey represents SCC groups, medium grey represents WCC groups, dark grey represents NCC groups.} \\label{fig:Khist} \\end{figure}  \\begin{figure*} \\centering \\includegraphics[scale=0.50,angle=-90]{Ktcool_3007.eps} \\caption{$K_{0}$ vs. CCT. Magenta is the best fit for all groups, blue for only SCC groups and red for WCC+NCC groups.} \\label{fig:entropy} \\end{figure*}  \\subsection{Radio properties}  \\subsubsection{CRS fractions-CC/NCC dichotomy}\\label{CRSfract}  Table~\\ref{CCT} also lists whether or not there is a central radio source present in a group. We see that while \\textit{all} NCC groups have a CRS, the fractions of SCCs and WCCs containing a CRS are 77\\% and 57\\%, respectively. The overall CRS fraction for CC groups is 70\\%. Figure~\\ref{fig:fraction} shows the CRS fractions in the group sample.  \\subsubsection{Total radio luminosity vs.~central cooling time}  \\cite{ 2009A&A...501..835M} show that there is an anti-correlation trend between the CCT and the total radio luminosity for CC clusters (correlation coefficient of $-0.63$), which breaks down for cooling times shorter than 1~Gyr.  Figure~\\ref{fig:Rtot} shows the same plot for groups. We do not find indications of a trend between the two quantities. Here we show the best fit obtained for the CC clusters from \\cite{ 2009A&A...501..835M} to highlight the difference between clusters and groups. The power-law fit for the CC clusters from \\cite{ 2009A&A...501..835M} is given by:  \\begin{equation} L_{\\mathrm{tot}} = (0.041\\pm0.016)\\times (t_{\\mathrm{cool}})^{-3.16\\pm0.38}~. \\end{equation}  It is interesting to note that \\textit{all} SCC groups and most of the WCC groups show a much lower radio output than the best fit for clusters. This was first alluded to in \\cite{ 2009A&A...501..835M}, where the groups in that sample (all SCC) were clear outliers and we confirm this for the first time with a large sample of groups. We discuss this in detail in Sect.~\\ref{AGNact}. \\begin{figure*} \\includegraphics[scale=0.35,angle=-90]{tcoolRtot_3007.eps} \\includegraphics[scale=0.35,angle=-90]{SMBHradfinal_3007.eps} \\caption{Left: Total radio luminosity vs. CCT with best-fit line for \\textit{clusters} from \\cite{ 2009A&A...501..835M}. Right: Total radio luminosity vs. mass of SMBH. The blue asterisks represent SCC groups, the green triangles represent WCC groups and the red circles represent NCC groups.} \\label{fig:Rtot} \\end{figure*}  \\subsubsection{Total radio luminosity vs.~SMBH mass}  The total radio luminosity shows no trend with the SMBH mass (Fig.~\\ref{fig:Rtot}). Classifying the sample as SCC, WCC, and NCC also does not yield any discernible correlations. This is in contrast with the HIFLUGCS sample, which shows a weak correlation for the SCC clusters (correlation coefficient of 0.46). We calculate a correlation coefficient of 0.29 for all groups and 0.20 for only SCC groups. In Table~\\ref{meanSMBHmassLrad} we present the mean SMBH masses and radio luminosities for the different CC types along with their standard errors to investigate whether different CC groups have systematically different masses and/or luminosities. We observe that the NCC groups have a systematically higher SMBH mass and radio luminosity than the SCC and WCC groups.  \\begin{table*} \\begin{center} \\caption{Mean SMBH masses and radio luminosities for different CC groups.}\\label{meanSMBHmassLrad} \\begin{tabular}{|c|c|c|} \\hline Group & Mean SMBH mass ($10^{9}$ solar masses) & Mean radio luminosity ($10^{42}$ erg/s) \\\\ \\hline SCC &1.17$\\pm$0.18&0.0050$\\pm$0.0023\\\\ WCC &1.39$\\pm$0.24&0.0034$\\pm$0.0028\\\\ NCC &1.98$\\pm$0.27&0.24$\\pm$0.15\\\\ \\hline \\end{tabular}  \\end{center}  \\end{table*}  \\subsection{BCG properties} We present the scaling relation between the BCG and the cluster/group X-ray luminosity and mass in Figs.~\\ref{fig:LxLbcg} and \\ref{fig:M500Lbcg}. The fits shown here are for a combined relation for groups and clusters. The details are summarized in Table~\\ref{scat}. The derivation of the scatter is explained in Appendix~\\ref{Scatcalc}. We observe that most group BCGs lie above the best-fit relations. Additionally, extending the scaling relations from clusters to groups leads to a higher intrinsic scatter in most cases. \\subsubsection{BCG luminosity vs.~X-ray luminosity}  Figure~\\ref{fig:LxLbcg} shows the relation between the two quantities for both the HIFLUGCS clusters and the groups, with fits for all the members, SCC members, and non-SCC members. The correlation coefficients are 0.65, 0.79, and 0.46 for all, SCC and non-SCC systems, respectively.  It is seen that most group BCGs have a higher luminosity than expected from the derived scaling relation. The best-fit power-law relation is given by \\begin{equation} \\left ( \\frac{L_{\\mathrm{BCG}}}{10^{11}L_{\\odot}} \\right )  = c\\times \\left ( \\frac{L_{\\mathrm{X}}}{10^{44} \\mathrm{ergs/s}} \\right )^{m}~, \\end{equation} where $m = 0.34\\pm0.03$ and $c = 6.98\\pm0.16$ for all systems, $m = 0.32\\pm0.03$ and $c = 6.89\\pm0.20$ for SCC systems, and $m = 0.43\\pm0.09$ and $c = 7.01\\pm0.30$ for non-SCC systems. We observe that when the relations are extended to the group regime, there is no significant difference in slopes and normalisations for the subsets. \\begin{figure*} \\centering \\includegraphics[scale=0.50, angle=270]{LxLbcg_3007.eps} \\caption{BCG luminosity vs.~X-ray luminosity. The black line is the best fit for all systems, the blue line for only SCC systems, and the red line for non-SCC systems.} \\label{fig:LxLbcg} \\end{figure*} \\begin{figure*} \\centering \\includegraphics[scale=0.50, angle=270]{M500Lbcg_3007.eps} \\caption{BCG luminosity vs.~$ M_{500}$. The black line is the best fit for all systems, the blue line for only SCC systems, and the red line for non-SCC systems.} \\label{fig:M500Lbcg} \\end{figure*}  \\subsubsection{BCG luminosity vs.~\\texorpdfstring{$ M_{500}$}{M500}}  Figure~\\ref{fig:M500Lbcg} shows that the luminosity of the BCG grows with the cluster/group mass. The correlation coefficients are 0.65, 0.86 and 0.44 for all, SCC and non-SCC systems respectively. The best fit powerlaw relation is given by: \\begin{equation} \\left ( \\frac{L_{\\mathrm{BCG}}}{10^{11}L_{\\odot}} \\right )  = c \\times \\left (\\frac{M_{500}}{10^{14} M_{\\odot}} \\right)^{m} \\end{equation} where $m = 0.49\\pm0.03$ and $c = 4.74\\pm0.12$ for all systems, $m = 0.52\\pm0.04$ and $c = 5.15\\pm0.14$ for SCC systems, and $m = 0.56\\pm0.08$ and $c = 4.17\\pm0.03$ for non-SCC systems. We see that the normalisations for the SCC systems are around 23\\% higher than those for non-SCC systems.  The total mass $M_{500}$ was calculated from the virial temperature (taken from \\citealt{ 2011A&A...532A.133M} and \\citealt{ Eckmiller} for the clusters and groups, respectively) through the scaling relation given in \\cite{2001A&A...368..749F}: \\begin{equation} \\left ( \\frac{M_{500}}{10^{13} h^{-1}_{71} M_{\\odot}}  \\right ) = (2.5\\pm0.2)\\left ( \\frac{kT_{\\mathrm{vir}}}{1 \\mathrm{keV}} \\right )^{1.676\\pm0.054}~. \\label{M500} \\end{equation}  \\begin{table*} \\centering \\caption{Best-fit results and scatter for BCG-cluster scaling relations: ``stat'' refers to statistical scatter and ``int'' refers to intrinsic scatter.} \\begin{tabular}[]{| c c  c  c  c  c  c |} \\hline \\textbf{Category} & \\textbf{Slope} & \\textbf{Normalisation} & $\\sigma_{\\mathrm{int},L_{\\mathrm{X}}}$ & $\\sigma_{\\mathrm{stat},L_{\\mathrm{X}}}$ & $\\sigma_{\\mathrm{int},L_{\\mathrm{BCG}}}$ & $\\sigma_{\\mathrm{stat},L_{\\mathrm{BCG}}}$ \\\\ \\hline $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters & 0.36$\\pm$0.03 &4.54$\\pm$0.34 & 0.67 & 0.07 & 0.24 & 0.03 \\\\ $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters+Groups &0.34$\\pm$0.03&6.98$\\pm$0.16& 0.85 & 0.08 & 0.21 & 0.03 \\\\ $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters (SCC) &0.32$\\pm$0.03&5.15$\\pm$0.38& 0.64 & 0.08 & 0.21 & 0.02 \\\\ $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters+Groups (SCC) &0.32$\\pm$0.03&6.89$\\pm$0.20& 0.57 & 0.08 & 0.18 & 0.02 \\\\ $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters (NSCC) &0.50$\\pm$0.07&3.49$\\pm$5.09&0.61 & 0.06 & 0.21 & 0.03 \\\\ $L_{\\mathrm{X}}-L_{\\mathrm{BCG}}$ Clusters+Groups (NSCC) &0.43$\\pm$0.09&7.01$\\pm$0.30& 0.64 & 0.07 & 0.29 & 0.03 \\\\ \\hline &&&$\\sigma_{\\mathrm{int},M_{500}}$ & $\\sigma_{\\mathrm{stat},M_{500}}$ & $\\sigma_{\\mathrm{int},L_{\\mathrm{BCG}}}$ & $\\sigma_{\\mathrm{stat},L_{\\mathrm{BCG}}}$ \\\\ \\hline $M_{500}-L_{\\mathrm{BCG}}$ Clusters &0.62$\\pm$0.05&3.52$\\pm$0.28& 0.30 & 0.07 & 0.19 & 0.04 \\\\ $M_{500}-L_{\\mathrm{BCG}}$ Clusters+Groups &0.49$\\pm$0.03&4.74$\\pm$0.12& 0.40 & 0.07 & 0.20 & 0.03 \\\\ $M_{500}-L_{\\mathrm{BCG}}$ Clusters (SCC) &0.62$\\pm$0.01&4.30$\\pm$0.29& 0.21 & 0.07 & 0.13 & 0.04 \\\\ $M_{500}-L_{\\mathrm{BCG}}$ Clusters+Groups (SCC) &0.52$\\pm$0.04&5.15$\\pm$0.14& 0.31 & 0.06 & 0.16 & 0.03 \\\\ $M_{500}-L_{\\mathrm{BCG}}$ Clusters (NSCC) &0.75$\\pm$0.09&2.55$\\pm$0.36& 0.31 & 0.07 & 0.23 & 0.05 \\\\ $M_{500}-L_{\\mathrm{BCG}}$ Clusters+Groups (NSCC) &0.56$\\pm$0.08&4.17$\\pm$0.03& 0.53 & 0.07 & 0.29 & 0.03 \\\\ \\hline \\end{tabular} \\label{scat} \\end{table*} ",
        "conclusions": "\\label{Summary}  With a sample of 26 Chandra galaxy groups we have peformed a study of the ICM cooling, AGN feedback and the BCG properties on the galaxy group scale. The major results of our study are as follows. \\begin{itemize} \\item The group sample has similar SCC, WCC, and NCC fractions to the HIFLUGCS cluster sample. \\item We find that 23\\% of the groups that have CCT $\\leq$ 1 Gyr do not show a central temperature drop. We speculate that this could be due to a partial reheating of the cool core in the past. Additionally, we also speculate that this might be a characteristic feature of fossil groups. \\item An increase in the CRS fraction with decreasing CCT is \\textit{not} seen, unlike for the HIFLUGCS sample. This is the first indication of differences between clusters and groups in the AGN heating/ICM cooling paradigm. \\item There is no correlation seen between the CCT and the integrated radio luminosity of the CRS. We notice that CRSs for the SCC groups, in particular, have a much lower radio luminosity than clusters. \\item We extend the scaling relations between $L_{\\mathrm{BCG}} $ and global cluster properties ($L_{\\mathrm{X}} $ and $M_{500}$) into the group regime. Most group BCGs have a BCG luminosity above the best fit and we think this may be due to a higher stellar mass content in group BCGs than in cluster BCGs, for a given $L_{\\mathrm{X}}$ and $M_{500}$. \\item We have speculated that star formation is a possible, effective answer to the fate of the cool gas in groups where, because of fueling star formation, not enough gas (low as it is) is being fed to the SMBH and hence the radio output of the CRSs is not as high. This could also explain why some SCC groups do not show a CRS. \\end{itemize} In conclusion, we have demonstrated differences between clusters and groups vis-a-vis ICM cooling, AGN feedback, and BCG properties. These results lend support to the idea that groups are not simply scaled-down versions of clusters. In the future, it would be interesting to study the impact of these processes on scaling relations for galaxy groups."
    },
    "1402/1402.4269.txt": {
        "abstract": "{We extend the resummation method of Anselmi \\& Pietroni (2012) to compute the total density power spectrum in models of quintessence characterized by a vanishing speed of sound. For standard $\\Lambda$CDM cosmologies, this resummation scheme allows predictions with an accuracy at the few percent level beyond the range of scales where acoustic oscillations are present, therefore comparable to other, common numerical tools. In addition, our theoretical approach indicates an approximate but valuable and simple relation between the power spectra for standard quintessence models and models where scalar field perturbations appear at all scales. This, in turn, provides an educated guess for the prediction of nonlinear growth in models with generic speed of sound, particularly valuable since no numerical results are yet available.} ",
        "introduction": "Ever since the discovery of the cosmic acceleration found by the observations of type Ia supernovae in 1998 \\cite{RiessEtal1998, PerlmutterEtal1999}, a tremendous amount of observational and theoretical effort has been put in place in order to understand the cause of this phenomenon.  A plethora of alternative models attribute the observed acceleration to either the presence of a mysterious energy component dubbed dark energy (DE) which dominates at the current epoch, or to  a modification of the gravitational laws on cosmological distances (see, e.g. \\cite{CopelandSamiTsujikawa2006} for an overview of proposed theoretical models). While the standard $\\Lambda$CDM model, where the acceleration is driven by a cosmological constant $\\Lambda$, provides a good fit to observations, several models cannot, so far, be ruled-out. The next generation of large-scale structure (LSS) surveys such as EUCLID \\cite{LaureijsEtal2011} and DESI\\footnote{\\href{http://desi.lbl.gov/}{http://desi.lbl.gov/}} will provide an unprecedented determination of observables such as the weak lensing shear and galaxy power spectra over a very large range of scales, and will be therefore sensitive to the details of structure formation, representing a unique opportunity to discriminate between competing models. However, the improvement in the quality and quantity of observational data represents a challenge to the accuracy of theoretical predictions, particularly for non-standard models.  A fundamental complication, in this respect, is given by the nonlinear nature of structure formation. Indeed many cosmological probes of matter density correlations will present small statistical errors on scales smaller than $\\sim 100\\Mpc$, where gravitational instability is responsible for relevant nonlinear corrections to linear theory predictions. The solution to the problem entails two goals: to understand and predict these corrections on one hand and, on the other, to achive a fast way (seconds) to compute numerically such predictions, in order to efficiently sample the theory parameter space in the data analysis. Our current understanding of the nonlinear regime, rather limited even for the standard paradigm, is very poor for non--standard cosmological models. In the following we will present an important step in this direction for the case of models with perturbations in the quintessence field.  Quintessence is one of the most popular models of DE, where the acceleration of the Universe is attributed to a scalar field with negative pressure \\cite{ZlatevWangSteinhardt1999}. Its standard formulation involves a minimally coupled canonical scalar field, whose fluctuations propagate at the speed of light $c_s=1$. In this case, sound waves maintain the quintessence homogeneous on scales smaller than the Hubble radius $H^{-1}$ \\cite{FerreiraJoyce1997}. The quintessence contribution to the total energy density affects the  expansion of the Universe and the growth of dark matter perturbations, allowing observational constraints on the DE equation of state $w(t)=p_Q(t)/\\rho_Q(t)$, where $p_Q$ and $\\rho_Q$ are the pressure and energy density of the scalar field. In this set-up, quintessence clustering effects take place only on scales of order $H^{-1}$, and their detection is therefore severely limited by cosmic variance. However, this may not be the case for more general quintessence models where the scalar field kinetic term is non-canonical\\footnote{ The name ``quintessence'' is usually reserved for theories where the dark energy is given by a scalar field with a canonical kinetic term, however, following previous works on the subject, we will refer as quintessence to a generic minimally coupled scalar field, with either canonical or non-canonical kinetic term.}.  In fact,  in \\cite{CreminelliEtal2006B,CreminelliEtal2009}, following the approach of \\cite{CheungEtal2008},  the authors constructed the most general effective field theory  describing quintessence perturbations around a given Friedman-Robertson-Walker background space-time, showing the existence of theoretically consistent models where the fluctuations propagate at a speed $c_s\\neq 1$.   In particular, by a careful analysis of the theoretical consistency of the models, namely, the classical stability of the perturbations and the absence of ghosts, ref.~\\cite{CreminelliEtal2009} concluded that the region of parameter space with $w<-1$ (the so- called ``phantom regime'') is in general  excluded, unless the  fluctuations are characterized by a practically  vanishing, imaginary speed of sound, e.g. $|c_s|\\lesssim 10^{-15}$ (note that observational limits  do not exclude  the case $w<-1$  \\cite{Planck2013parameters,RestEtal2013}). In the latter case,   the stability of the perturbations would be guaranteed by the presence of higher derivative operators with negligible effects on cosmological scales. As was emphasized in  \\cite{CreminelliEtal2006B,CreminelliEtal2009}, such a tiny speed of sound should not be considered as a fine tuning, since in the limit $c_s\\to 0$ one recovers the  ghost condensate model of ref.~\\cite{ArkaniHamedEtal2004A}, which is invariant under a shift symmetry of the field,  and hence  a value of $c_s\\simeq 0$ can be interpreted as a deformation of the limit where such symmetry is recovered.  What makes these models with $c_s\\simeq 0$ of particular interest is that quintessence perturbations can cluster on all observable scales, opening up to a rich phenomenology. In fact, for such models the quintessence field is expected to affect not only the growth history of dark matter through the background evolution but also by actively participating to structure formation, both at the linear and nonlinear level. In the linear regime, the observational consequences of clustering quintessence have been investigated by several authors: see for instance \\cite{WellerLewis2003, BeanDore2004, DeDeoCaldwellSteinhardt2003, Hannestad2005, EricksonEtal2002, DePutterHutererLinder2010, SaponeKunz2009} for large scale CMB anisotropies, \\cite{Takada2006} for galaxy redshift surveys, \\cite{TorresRodriguezCress2007} for neutral hydrogen observations or \\cite{HuScranton2004, CorasanitiGiannantonioMelchiorri2005} for the cross-correlation of the Integrated Sachs-Wolfe effect in the CMB and the LSS. However, the possibility to detect quintessence perturbations effects focusing on observations in the linear regime alone is limited by large degeneracies with cosmological parameters, particularly when quintessence clusters on all relevant scales. In this case, in fact, such effects consist essentially in a change of the fluctuations amplitude. On the other hand the relation between the linear {\\em and} nonlinear growth of density perturbations can be affected by clustering quintessence in a peculiar way, leading to specific signatures that could in principle allow to distinguish these models from $\\Lambda$CDM and homogeneous quintessence cosmologies \\cite{SefusattiVernizzi2011, DAmicoSefusatti2011, AnselmiBallesterosPietroni2011}.  Robust predictions of the nonlinear dynamics of the structure formation process ultimately require numerical simulations. The case of homogeneous quintessence (i.e., $c_s=1$), as it affects only the background evolution has been investigated by means of N-body simulations in several articles (see for instance \\cite{Baldi2012, KuhlenVogelsbergerAngulo2012} and references therein). The case of clustering quintessence, however, represents a much more difficult problem and numerical results are still lacking. The extension of common fitting formulas such as \\texttt{halofit} \\cite{SmithEtal2003, TakahashiEtal2012, Zhao2013}, based in turn on numerical simulations, is also, {\\em a priori} not possible. Precisely for this reason, the study of analytical approximations in cosmological perturbation theory (PT) to investigate the mildly nonlinear regime is particularly relevant. On one hand it can give insights in the physical processes at play. One of the main results of this work, in fact, consists in showing how fitting functions valid for tested, standard cosmologies can be used in clustering quintessence models. Such extensions have been already considered in the literature without a proper theoretical justification \\cite{SaponeKunz2009, AmendolaKunzSapone2008, SaponeKunzAmendola2010}. On the other hand, PT predictions can as well provide useful checks for future numerical results, which are likely to comprise significant approximations in their description of the dynamics of quintessence perturbations.  The nonlinear regime of structure formation in quintessence models with $c_s=0$ has been studied in previous literature. Several works focussed on the spherical collapse of structures in the presence of quintessence perturbations \\cite{MotaVanDeBruck2004, NunesMota2006, AbramoEtal2007}. In addition, ref.~\\cite{CreminelliEtal2010} makes use of the spherical collapse model \\cite{GunnGott1972} (see \\cite{BasseEggersBjaeldeWong2011, BatistaPace2013} for the case of arbitrary sound speed) along with the Press-Schechter formalism \\cite{PressSchechter1974}, to provide predictions for the halo mass function. We will consider here, instead, analytical predictions in cosmological perturbation theory for the mildly nonlinear regime of the density power spectrum (see \\cite{BernardeauEtal2002} for a review of standard PT).  In the context of standard, $\\Lambda$CDM cosmologies, perturbative techniques have witnessed a significant progress, motivated by the need for accurate predictions for Baryon Acoustic Oscillations measurements, following the seminal papers of Crocce \\& Scoccimarro \\cite{CrocceScoccimarro2006A, CrocceScoccimarro2006B, CrocceScoccimarro2008}. Several different \u00d2resummed\u00d3 perturbative schemes and approaches have been proposed in the last few years \\cite{MatarresePietroni2007, TaruyaHiramatsu2008, Pietroni2008, BernardeauCrocceScoccimarro2008, BernardeauVanDeRijtVernizzi2012, JurgensBartelmann2012, BlasGarnyKonstandin2014}, along with public codes implementing some of the methods \\cite{CrocceScoccimarroBernardeau2012,TaruyaEtal2012}. The key improvement over standard PT consisted in reorganizing the perturbative series in a more efficient way by means of a new building block, the nonlinear propagator, a quantity measuring the response of the density and velocity perturbations to their initial conditions. The possibility to analytically compute this new object revolutionized indeed the cosmological perturbation theory approach.  A first extension of standard PT to the specific case of clustering quintessence can be found in \\cite{SefusattiVernizzi2011}, where the pressure-less perfect fluid equations have been worked out, while in \\cite{DAmicoSefusatti2011, AnselmiBallesterosPietroni2011} it was shown how to compute the nonlinear power spectrum (PS) by means of the Time-Renormalization Group (TRG) method of \\cite{DAmicoSefusatti2011, AnselmiBallesterosPietroni2011}. However, the computation of the nonlinear propagator for these cosmologies was still lacking. This is the first achievement of the present work, allowing the extension of a large set of powerful resummed techniques to clustering Dark Energy models.  We compute the nonlinear PS exploiting the resummation scheme proposed by Anselmi \\& Pietroni \\cite{AnselmiPietroni2012} (hereafter AP). Such scheme takes into account the relevant diagrams in the perturbative expansion and proposes a new interpolation procedure between small and large scales. What makes AP so appealing is the possibility to extend the predictions, for the first time, to scales beyond the Baryonic Acoustic Oscillation (BAO) range, together with a fast and easy code implementation. It is important to stress that the AP approach, as other perturbative techniques, is intrinsic limited by the emergence of multi-streaming effects  at small scales, i.e. the generation of velocity dispersion due to orbit crossing. Clearly, the breaking-down of the pressureless perfect fluid (PPF) approximation employed as a starting point for the fluid equations has nothing to do with the perturbative scheme used. The AP prediction agrees with N-body simulations at the 2\\% level accuracy on BAO scales while at smaller scales shows discrepancies only of a few percents up to $k \\simeq 1\\kMpc$ at $z \\gtrsim 0.5$. As we expect, the level of the agreement is consistent with the findings of \\cite{PueblasScoccimarro2009} where the authors estimated the departures from the PPF assumption by means of high resolution N-body simulations. Clearly, it would be desirable to extend our results beyond the PPF description. However, so far, even in the standard scenario this is still subject of intense investigations (see \\cite{PueblasScoccimarro2009, ValageasNishimichi2011A, PietroniEtal2011, CarrascoHertzbergSenatore2012, PajerZaldarriaga2013, ValageasNishimichiTaruya2013} for an incomplete list of contributions).  Taking advantage of the AP approach, we will show that additional nonlinear corrections due DE clustering become larger than $1\\%$ for scales smaller than BAO for 10\\% variation from $w=-1$. Thus it allows us to finally test the approximation employed in the TRG predictions of \\cite{AnselmiBallesterosPietroni2011} and the assumptions made in several works forecasting the detectability of dark energy models in presence of quintessence perturbations \\cite{SaponeKunz2009, AmendolaKunzSapone2008, SaponeKunzAmendola2010}. Furthermore, as a by-product, we provide a theoretically motivated mapping from the nonlinear matter power spectrum in {\\em smooth} quintessence models to the nonlinear, total density power spectrum in {\\em clustering} quintessence models that could possibly serve as a consistency check for numerical results once these will become available. This is not of small significance, given the challenge posed by accurate simulations of structure formation in this kind of models.  This paper is organized as follows. In Section \\ref{EqMotionLinear} we review the equations of motion for the quintessence--dark matter fluid. In Section \\ref{dynamics} we recast the starting fluid equations in a compact form and make explicit the nonlinear behavior of perturbations. In Section \\ref{StatisticalObs}, in order to compute the nonlinear density power spectrum, we apply the selected resummation scheme to the clustering quintessence fluid. In Section \\ref{results} we present our results and a well motivated mapping from the smooth to the clustering quintessence PS. Section \\ref{discussion} is devoted to our conclusions and possible future developments. Some details about the resummation scheme and a comparison with recent numerical simulations for a $\\Lambda$CDM cosmology are included as Appendices.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{discussion}  Regardless any proper theoretical motivation or degree of naturalness, quintessence models characterized by a vanishing speed of sound are extremely interesting under a phenomenological point of view. In this case, in fact, quintessence perturbations are present at all scales and, more importantly, in the single-stream approximation, they are comoving with perturbations in the matter component, i.e.~they share the same velocity field. This allows to easily extend sophisticated techniques in cosmological perturbation theory already developed for standard, $\\Lambda$CDM cosmologies, \\cite{SefusattiVernizzi2011}.  At the same time, $c_s=0$ represents a limiting case of a much broader class of quintessence models,  parametrized by the equation of state parameter $w$ (with allowed values close to $-1$) and the  sound speed parameter $c_s$ taking values, in principle, in the relatively wide interval $0\\le c_s\\le1$. Needless to say, the possible detection of any departure from the standard values $w=-1$ and $c_s=1$ would be of great significance toward an understanding of the mechanism driving the observed acceleration of the Universe. Still, the theoretical description of the nonlinear evolution of quintessence perturbation poses a remarkable challenge. The difficulty of the task is highlighted by the fact that no numerical simulation for these quintessence models has been performed to date. Furthermore, it is easy to imagine that when they {\\em will} be performed they will likely involve relevant approximations. Precisely for this reason any analytical insight into the nonlinear behavior of a central quantity as the density power spectrum is greatly valuable. Approximate solutions for the power spectrum of (standard) models with $c_s=1$ and (extreme) models with a vanishing speed of sound will therefore mark the limits of the broader set of solutions for models with intermediate values of $c_s$ and provide boundaries for useful interpolations, both under a practical as under a theoretical point of view.  With this goal in mind, we applied, in this paper, the resummation scheme proposed in \\cite{AnselmiPietroni2012} to quintessence models with $c_s=0$, providing predictions for the total (i.e. matter and quintessence) density power spectrum, in the standard single-fluid approximation. This approach extends the range of validity of perturbative solutions beyond the acoustic oscillations scales up to $k\\sim 0.6\\kMpc$ with an accuracy below the 2\\% level at redshifts $z\\ge 0.5$, therefore comparable to the accuracy delivered by the  \\texttt{halofit} fitting formula \\cite{SmithEtal2003, TakahashiEtal2012} or the emulator approach of \\cite{HeitmannEtal2014}.  In addition, our analytical investigation suggests an approximate but useful and simple mapping between the nonlinear power spectra of models with $c_s=1$ and models with $c_s=0$ when all other cosmological parameters are held at the same values. We show, in fact, that the ratio between these two quantities is very close to the ratio of the square of their respective linear growth factors, eq.~(\\ref{nlApp}). This is the likely consequence of the fact, already suggested in earlier works \\cite{SefusattiVernizzi2011, AnselmiBallesterosPietroni2011, DAmicoSefusatti2011}, that nonlinear growth is driven predominantly by the matter component.  This is quite an interesting result as it provides, for the first time, a theoretical motivation for simple predictions of the density power spectrum that can make use of tools as  \\texttt{halofit} or more in general results from existing, numerical simulations of more standard models. Waiting for a proper N-body description of clustering quintessence models, our findings will enable more robust assessments of the ability of future cosmological probes to constrain this class of models. We will address these issues in future work."
    },
    "1402/1402.0603_arXiv.txt": {
        "abstract": "We report on spatially resolved X-ray spectroscopy of the north-eastern part of the mixed morphology supernova remnant (SNR) W28 with {\\it XMM-Newton}. The observed field of view includes a prominent and twisted shell emission forming the edge of this SNR as well as part of the center-filled X-ray emission brightening toward the south-west edge of the field of view.  The shell region spectra are in general represented by an optically thin thermal plasma emission in collisional ionization equilibrium with a temperature of $\\sim$0.3~keV and a density of $\\sim$10~cm$^{-3}$, which is much higher than the density obtained for inner parts. In contrast, we detected no significant X-ray flux from one of the TeV $\\gamma$-ray peaks with an upper-limit flux of 2.1$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ in the 2--10~keV band. The large flux ratio of TeV to X-ray, larger than 16, and the spatial coincidence of the molecular cloud and the TeV $\\gamma$-ray emission site indicate that the TeV $\\gamma$-ray of W28 is $\\pi^{0}$-decay emission originating from collisions between accelerated protons and molecular cloud protons. Comparing the spectrum in the TeV band and the X-ray upper limit, we obtained a weak upper limit on the magnetic field strength $B\\lesssim$ 1500~$\\mu$G. ",
        "introduction": "Supernova remnants (SNRs) are one of the most promising acceleration sites of cosmic rays up to $\\sim10^{15.5}$~eV~(the knee energy). \\citet{1995Natur.378..255K} discovered synchrotron X--rays from the shell of SN 1006, indicating the existence of extremely high-energy electrons up to $\\sim$~TeV produced by the first-order Fermi acceleration. Following this discovery, the synchrotron X-ray emission has been discovered from a few more young shell-type SNRs, such as RX~J1713.7--3946 \\citep{1997PASJ...49L...7K}, RCW~86 \\citep{2000PASJ...52.1157B,2001ApJ...550..334B}, and RX~J0852.0$-$4622 \\citep{2000PASJ...52..887T}. On the other hand, TeV $\\gamma$--rays have also been detected from some SNRs \\citep{2004Natur.432...75A,2006A&A...448L..43A}. The radiation of TeV $\\gamma$-ray is explained by (1) Inverse-Compton scattering (IC) of cosmic microwave background photons by the same high energy electron giving rise to the X-ray synchrotron emission, (2) non-thermal bremsstrahlung by high energy elections or (3) the decay of neutral pions that are generated by collisions between high energy protons and dense interstellar matter.  Most of the SNRs with such a X-ray and TeV $\\gamma$-ray evidence have been young with an age of less than several thousands of years. Recently, Fermi discovered GeV $\\gamma$-rays from several SNRs, such as W44 \\citep{2010Sci...327.1103A}, IC443 \\citep{2010ApJ...712..459A}, W51C \\citep{2009ApJ...706L...1A}, and so on. Interesting fact is that they are not only young but rather old SNRs compared with the TeV $\\gamma$-ray emitting SNRs. It can be due to the escape of high energy particles from the shocks \\citep{2011MNRAS.410.1577O,2011ApJ...731...87E,% 2012MNRAS.421..935L,2012A&A...541A.153T,2013MNRAS.429.1643N}. Most SNRs with GeV $\\gamma$-rays are interacting with molecular clouds, so it may be related with the escape. However, it is still unknown the detailed picture how particles escapes from the shocks.   The SNR W28, locating at ($l$, $b$) = ($\\timeform{6D.4}$, $-\\timeform{0D.1}$), is an interesting target as a cosmic-ray accelerator from which both GeV and TeV $\\gamma$-rays were detected from the eastern edge of the radio shell~(\\cite{2009ApJS..183...46A,2008A&A...481..401A}). The diameter and the distance to W28 are 48 arcmin and 1.9~kpc, respectively~\\citep{2002AJ....124.2145V}. The age seems to be several times 10$^{4}$ year, which means the remnant is middle-aged~\\citep{2002ApJ...575..201R}. W28 is classified as ``mixed--morphology'' SNR, showing center-filled X-rays and a shell-like radio emission. The shell-like radio emission peaks at the northern and northeastern boundaries where interaction of the SNR matter with the molecular cloud is established~\\citep{1981ApJ...245..105W}. This interaction was revealed by a lot of OH maser spots \\citep{1994ApJ...424L.111F,1997ApJ...489..143C,2005ApJ...620..257H}, which are signposts of molecular interactions and the location of high density shocked gas ($n \\ge$10$^{4}$ cm$^{-3}$; \\cite{1999PASJ...51L...7A}). TeV $\\gamma$-ray emission was detected by H.E.S.S.~\\citep{2008A&A...481..401A} from the eastern edge of the radio shell. Recently, {\\it{Fermi}} and {\\it AGILE} also detected $\\gamma$-rays at 100~MeV $\\textless E \\textless$ 100~GeV ~\\citep{2010ApJ...718..348A,2010A&A...516L..11G} from the same region of the TeV emission. In contrast, the emission in the X-ray band had been believed to comprise of thermal radiation. \\citet{2002ApJ...575..201R} evaluated temperatures of 1.5~keV in the southwest, 0.56~keV in the northeast, and the central region requires two-temperature plasma with 0.6~keV and 1.8~keV. The long ionization timescales in the northeast and central region imply that the gas is close to the ionization equilibrium. Recently, Suzaku discovered that the thermal X-ray emission from the inner region is over-ionized, implying that the plasma may underwent the sudden rarefaction \\citep{2012PASJ...64...81S}. These facts imply that W28 is an ideal target to study the particle escape from the shock.  We analyzed XMM-Newton archival data of the north--eastern part of W28, where the molecular clouds, OH maser spots, bright X-ray shells, GeV and TeV emission were detected. XMM-Newton has large effective area and high angular resolution. These characteristics enable us to carry out high quality spatially resolved spectroscopy. In \\S~2, we present the observation log and data reduction method. Imaging and spectral analyses are shown in \\S~3 and \\S~4, respectively. From one of the TeV $\\gamma$-ray emission region, we obtained only the upper limit on the X-ray flux. Discussions are made on the basis of these results in \\S~5 on the nature of the thermal and non-thermal components in multi-wavelength. Finally we summarize our results in \\S~6. ",
        "conclusions": "We analyzed the XMM-Newton data of the north-eastern part of the supernova remnant W28. The observed X-ray image showed the bright and twisted north-eastern shell and the inner emission region which is part of the center-filled emission brightening toward the south-west end of the field of view.  The X-ray emission from north-eastern shell is found to reach collisional ionization equilibrium state and can be fitted well with a single temperature optically thin thermal emission model with $kT$ of $\\simeq$0.3~keV. From the emission measure and the apparent volume, the electron density is found to be as high as $\\simeq$10~cm$^{-3}$. Since a bunch of molecular cloud spatially coincides with the outer edge of a part of the shell, this high density is due to the collision of the plasma with the molecular cloud.  In contrast, there is no significant X-ray emission from one of the TeV $\\gamma$-ray peaks. We only obtained the 90\\% upper limit flux of 2.1$\\times$ 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ in the 2--10~keV band, assuming a power-law spectrum with the same photon index as in the TeV $\\gamma$-ray (= 2.66). The spatial coincidence of the molecular cloud and the TeV $\\gamma$-ray emission site suggests that TeV $\\gamma$-ray is hadronic origin. We calculate the spectra of hadrons including pions, kaons, nucleons and so on produced through proton-proton scatterings and their daughter particles (gamma-rays, electrons, neutrinos and so on), and found that $\\pi^{0}$-decay emission is dominant in the TeV $\\gamma$-ray band. A weak upper limit on the magnetic field strength is obtained $B\\lesssim$ 1500$\\mu$G from the X-ray flux upper limit."
    },
    "1402/1402.4436_arXiv.txt": {
        "abstract": "{We report the discovery of large--scale diffuse radio emission South-East of the galaxy cluster MACS\\,J0520.7-1328, detected through high sensitivity Giant Metrewave Radio Telescope 323~MHz observations. This emission is dominated by an elongated diffuse radio source and surrounded by other features of lower surface brightness. Patches of these faint sources are marginally detected in a 1.4 GHz image obtained through a re-analysis of archival NVSS data. Interestingly, the elongated radio source coincides with a previously unclassified extended X-ray source. We perform a multi--wavelength analysis based on archival infrared, optical and X-ray \\chandra data. We find that this source is a low--temperature ($\\sim$3.6 keV) cluster of galaxies, with indications of a disturbed dynamical state, located at a redshift that is consistent with the one of the main galaxy cluster MACS\\,J0520.7-132 (z=0.336). We suggest that the diffuse radio emission is associated with the non-thermal components in the intracluster and intergalactic medium in and around the newly detected cluster. We are planning deeper multi--wavelength and multi-frequency radio observations to accurately investigate the dynamical scenario of the two clusters and to address more precisely the nature of the complex radio emission. } ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=0.95\\textwidth]{fig1.pdf} \\caption{ {\\em Left panel:} 323~MHz GMRT full--resolution (11.8\\secpoint $\\times$ 7.8\\secpoint, p.a. 16.9\\degpoint) image of the 19\\minpoint$\\times$19\\minpoint area around the cluster M0520. Contours are spaced by a factor of two and start at $\\pm$3$\\sigma$=0.33 mJy/b (negative contours are dashed). The  dashed circle indicates the central region of 1 Mpc radius around the cluster M\\,0520 . {\\em Right panel:} zoom in of the area around the diffuse elongated source; radio contours (same as  in {\\em Left panel}) are overlaid on the DSS2 optical red image. The black circles show position of the three galaxies from DSS catalog, which are likely optical counterparts of radio sources S1 and S2. } \\label{fig:fig1} \\end{figure*}    \\begin{figure*} \\centering \\includegraphics[width=1.\\textwidth]{fig2.pdf} \\caption { {\\em Left panel}: 323~MHz GMRT low--resolution (21.1\\secpoint $\\times$ 17.2\\secpoint, p.a. 28.8\\degpoint) image of the same 19\\minpoint$\\times$19\\minpoint region around the cluster MACS\\,J0520 (same as in Fig. \\ref{fig:fig1}); black contours are spaced by a factor of two and start at $\\pm$3$\\sigma_{323\\rm MHz}$=0.82 mJy/b (negative contours are dashed). To highlight the position of discrete compact radio sources in the field, contours from the full--resolution image (see Fig. \\ref{fig:fig1}) are over--plotted in red, starting from $\\pm6\\sigma$=0.66 mJy/b. Dotted ellipses mark the location of diffuse emission, labelled D1 to D4. {\\em Right panel:}  \\vla 1.4 GHz contours (in blue) of an image obtained from re--processed NVSS data (45.8\\secpoint $\\times$ 29.9\\secpoint, p.a. -70.0\\degpoint); contours start at $\\pm2\\sigma_{1400\\rm MHz}=0.5$ mJy/b and are spaced by a factor of 2  (blue dashed are negative). The contours are overlaid on a 323~MHz GMRT image of similar resolution (41.5\\secpoint$\\times$33.5\\secpoint, p.a. 21.3\\degpoint), shown in grey scale (the grey dashed contour corresponds to 3$\\sigma_{323\\rm MHz}$=1.65  mJy/b). } \\label{fig:fig_2} \\end{figure*}   In the currently supported hierarchical scenario of structure formation, galaxy clusters assemble and evolve through mergers and accretion of smaller units of matter, at the intersection of cosmic filaments connecting clusters into the  large--scale structure of the Universe.  In addition to the most common methods for clusters detections -- based on the study of galaxy over-densities, lensing, extended X-ray emission due to thermal bremmstrahlung from the intracluster medium (ICM), and the sub-mm signal due to Inverse Compton scattering of CMB photons by hot intracluster electrons  (SZ effect) -- observations in the radio band can also probe the existence of clusters and/or large--scale structures. Powerful high-redshift radio galaxies are known to be efficient tools for detecting galaxy clusters at high redshifts (z$>$1.5; \\eg \\citealt{galametz13}, and reference therein). At lower redshift, tailed radio galaxies are being extensively and successfully used for cluster searches. Since they are known to be associated with galaxy clusters, they can be used to trace the high density environments in the Universe (\\eg \\citealt{gv09}). Several previously unobserved clusters have been actually identified thanks to the detection of tailed sources (\\eg  \\citealt{blanton03}; \\citealt{smolcic07}; \\citealt{gv09}; \\citealt{kantharia09}; \\citealt{mao2010}). \\\\ Beyond tailed radio galaxies, it is nowadays known that a fraction of merging clusters host diffuse Mpc--scale radio emission (such as radio halos and relics, that are named according with their location and observational properties; \\eg  see \\citealt{ferrari08}). This synchrotron emission is unrelated to individual cluster radio galaxies and directly probes the existence of non-thermal components (relativistic particles and magnetic fields) mixed with the thermal ICM. To date, all these sources are found to be associated with dynamically disturbed / merging systems (\\eg \\citealt{cassano10a}; \\citealt{cassano13}; \\citealt{feretti12}, for a recent review). The detection of such sources can become -- in principle -- a powerful tool to trace previously unidentified clusters. \\\\ Upcoming deep and wide-field radio surveys with new radio interferometers like LOFAR  (Low Frequency Array; \\eg \\citealt{vanhaarlem13}), ASKAP  (Australian Square Kilometer Array Pathfinder; \\eg \\citealt{johnston08}), MeerKAT (\\eg \\citealt{booth12}) and MWA   (Murchison Widefield Array; \\eg \\citealt{tingay2013}) will allow systematic searches for new galaxy clusters based on the detection of tailed radio galaxies and diffuse extended emission (see \\eg \\citealt{ensslin02}; \\citealt{cassano10b}). For these studies, LOFAR and MWA are particularly suitable, thanks to the very large field of view covered at low radio frequencies that makes them a powerful survey telescope (see \\ie the ongoing all-sky surveys LOFAR MSSS, \\citealt{heald13}, and MWA GLEAM, \\citealt{lss}, in preparation).\\\\ In this paper we present an example of a new cluster detection made possible thanks to deep low--frequency radio observations. We report the serendipitous discovery with the Giant Metrewave Radio Telescope (\\gmrt) at 323 MHz of new large--scale diffuse radio emission in the proximity of the cluster MACS\\,J0520.7-1328 (M\\,0520 hereinafter). \\\\ M\\,0520 is one of the X-ray brightest clusters of galaxies from the MAssive Cluster Survey (MACS, \\citealt{Ebeling10}). It is located at $z$=0.336 and has an X-ray luminosity of L$_{\\rm r500, [0.1-2.4 keV]} = 7.8 \\times 10^{44}$ erg/s \\citep{Mantz10}\\footnote{Corrected for our slightly different choice of cosmological parameters.}. M0520 is classified as a relatively relaxed system on the basis of its morphological properties, \\ie a good alignment between the galaxy and ICM distributions, and quite concentric X-ray surface brightness contours in its central 1.5 Mpc region \\citep[morphology code 2 in][]{Ebeling10}. This cluster is also found to be a lensing system, though no mass estimate has been derived from lensing studies \\citep{Horesh10}. M\\,0520 is identified in the Planck SZ cluster catalogue as PSZ1 G215.29-26.09,  with a mass estimated from the SZ of M$_{500}\\sim$6.15$\\times$10$^{14}$ M$_{\\odot}$ \\citep{plk2013}. \\\\ For our analysis, we made use of our new \\gmrt\\ radio observations and of multi--wavelength data available from archives and the literature. The paper is presented as follows:  in Sect. 2 we describe the new \\gmrt observations and data reduction, along with the analysis of the detected emission; in Sect. 3, archival \\chandra X-ray data analysis of the cluster is presented; in Sect. 4, we report optical/IR analysis based on available public catalogues. A multi--wavelength comparison is reported in Sect. 5, and the results are discussed in Sect. 6. \\\\ The adopted cosmology is  $\\Lambda$CDM, with ${\\rm H}_0$=71 km ${\\rm s}^{-1} {\\rm Mpc}^{-1}$, $\\Omega_{\\rm m} = 0.27$, $\\Omega_{\\Lambda} = 0.73$. At the redshift of the cluster 1$^{\\prime}$ corresponds to $\\sim$290 kpc. ",
        "conclusions": "We have reported high sensitivity \\gmrt\\ radio observations at 323 MHz of the galaxy cluster M\\,0520, that allowed the detection of complex diffuse radio emission. The main feature of this emission is located at a projected distance of $\\sim$8' SE of the cluster centre and was found to coincide with an extended X-ray source seen in the archival \\chandra image. We have performed a multi wavelength analysis based on archival X-ray and optical/IR data, that allowed us to classify it as a galaxy cluster, namely 1WGA\\,0521.  Our X-ray analysis allowed us to characterise this newly discovered cluster as a low--temperature, relatively small system. From the spectral fitting we derived a redshift that is consistent with the optical redshift of M\\,0520, z=0.336. However the uncertainties on our X-ray estimate (Sec. \\ref{sec:Xspec}, Table \\ref{tab:x}) do not allow us to firmly conclude that the two clusters lie at the same distance \\footnote{ Note that, at z$\\sim$0.336, the difference of 0.004 between the X-ray redshift estimates (see Table \\ref{tab:x}) corresponds to 25 Mpc in luminosity distance. If we also consider their uncertainties from spectral fitting, the distance between the two clusters could be much larger. }. Despite the fact that at z=0.336  their projected separation of $\\sim$2.3 Mpc would easily include the clusters virial radii, the  \\chandra\\ image does not show evidence of a bridge of X-ray emission between M\\,0520 and 1WGA0521, that could be expected in case of interaction (though this may be also due to the low sensitivity of the exposure). \\\\ On the other hand, our optical/IR analysis supports the hypothesis that the two clusters lie at similar redshifts, since their galaxy populations lie in statistically similar regions of the colour magnitude diagram. Optical spectroscopic observations and deeper X-ray observations are needed to accurately determine the 1WGA\\,0521 redshift and to investigate the dynamical scenario of these two systems and possible ongoing physical interaction. \\\\ The optical iso-density map derived from candidate cluster members shows that 1WGA\\,0521 has an elongated morphology, suggesting a disturbed dynamical state. Conversely, M\\,0520 seems more relaxed (as found by \\citealt{Ebeling10}), with a more regular shape, except for its inner region where the galaxy distribution is slightly elongated.   Our radio images at low resolution reveal a complex distribution of diffuse emission, that we have analysed by identifying four structures (namely D1, D2, D3 and D4, see Sect. \\ref{sec:radio}). The diffuse elongated source D1 is the brightest of these features. It is characterised by large--scale and  low--surface--brightness, and it is coincident with the newly detected, probably unrelaxed cluster of galaxy 1WGA\\,0521. These properties are typical for cluster--scale diffuse radio sources. Due to its projected location in the central region of the cluster, the emission could be a cluster radio halo with an elongated morphology. Such shape may also suggest the source is a radio relic (though located closer to the cluster centre with respect to what typically observed). The emission in this case would be associated with the passage of a shock wave, that may have either accelerated particles from the thermal ICM, or have revived fossil radio plasma from a previous activity of a radio--loud AGN (that could be one of the two nearby radio galaxies S1 and S2). Following the definition by \\cite{kempner2004}, the source would be classified as \\textit{radio gischt} in the former scenario, or a \\textit{radio phoenix} in the latter. The relic could be associated with the cluster M\\,520, rather than to 1WGA\\,0521. This hypothesis seems less likely, due to its large distance from the cluster centre ($\\sim$ 2.3 Mpc), much larger than what is typically observed, and the fact that M\\,520 seems to be a relaxed system, while relics are always found to be associated to merging clusters. Another possibility that cannot be firmly excluded is that D1 is actually the tail of a tailed radio galaxy, whose nucleus may be identified with S1 or S2. Our present radio data, however, do not allow us to disentangle between these possibilities; spectral and polarisation information is needed to clearly classify this source. \\\\ Beyond the prominent source D1, additional diffuse emission on large--scale is located around it (D2, D3, D4). The two elongated structures D2 and D3 are difficult to classify. They have an even lower brightness, their shape is filamentary, and they are located in a region in between 1WGA0521 and M\\,0520. D4 is even harder to interpret, due to its roundish shape and its location in a region where our present data does not show any significant optical or X-ray emission. We suggest that these features might be related to non-thermal components in the faint inter-galactic medium surrounding 1WGA0521, that is not detected by the current shallow X-ray data due to sensitivity limits, but may be revealed through more extensive observations (\\eg \\citealt{plck2013} and references therein). \\\\ Multi--frequency high sensitivity radio observations are essential to properly classify all the features of the complex extended radio emission, and to understand their origin. Follow up observations are planned. \\\\  Regardless the nature of the diffuse radio emission and its possible relation to the cluster's dynamical state, the most important result of our analysis is that our detection of large--scale diffuse radio emission has led us to the discovery of a new massive galaxy cluster, which was not previously known. Indeed, despite the fact that M0520 is a system in well studied cluster samples \\citep{Ebeling10,Mantz10,Horesh10} and that our X-ray and optical/IR analyses are based on archival data, the existence of the neighbouring cluster 1WGA0521 has been ignored up to now. Although it is not uncommon that X-ray surveys miss clusters (see \\eg \\citealt{cagnoni01}), in this case the radio observations have been crucial to spot its existence. \\\\ Thanks to on-going and future deep all-sky radio surveys (e.g. LOFAR--Surveys, EMU--ASKAP, MeerKAT--MighTEE, MWA GLEAM), diffuse radio sources will become a powerful tool for the discovery of dynamically disturbed galaxy clusters and large--scale structure merging/accretion events."
    },
    "1402/1402.6433_arXiv.txt": {
        "abstract": "Many dark energy models fail to pass the cosmic age test. In this paper, we investigate the cosmic age problem associated with nine extremely old Global Clusters (GCs) and the old quasar APM $08279+5255$ in the $R_h=ct$ Universe. The age data of these oldest GCs in M31 is acquired from the Beijing-Arizona-Taiwan-Connecticut system with up-to-date theoretical synthesis models. They have not been used to test the cosmic age problem in the $R_h=ct$ Universe in previous literature. By evaluating the age of the $R_h=ct$ Universe with the observational constraints from the type Ia supernovae and Hubble parameter, we find that the $R_h=ct$ Universe can accommodate five GCs and the quasar APM 08279+5255 at redshift $z=3.91$. But for other models, such as $\\Lambda$CDM, interacting dark energy model, generalized Chaplygin gas model and holographic dark energy model, can not accommodate all GCs and the quasar APM 08279+5255. It is worthwhile to note that the age estimates of some GCs are controversial. So, unlike other cosmological models, the $R_h=ct$ Universe can marginally solve the cosmic age problem, especially at high redshift. ",
        "introduction": "Many astronomical observations, such as type Ia supernovae (SNe Ia) \\cite{Riess,Perlmutter,Schmidt}, the cosmic microwave background (CMB) \\cite{Spergel,Komatsu,Planck}, gamma-ray bursts \\cite{Wang11a,Wang11b} and large-scale structure (LSS) \\cite{Tegmark}, indicate that the Universe is undergoing an accelerated expansion, which suggests that our universe may have an extra component like dark energy. The nature of dark energy is still unknown, but the simplest and most interesting candidate is the cosmological constant \\cite{Carroll}. This model can consist with most of astronomical observations. The latest observation gives that the present cosmic age is about $t_0=13.82$ Gyr in $\\Lambda$CDM model \\cite{Planck}, but it still suffer from the cosmic age problem \\cite{Yang,Wang}. The cosmic age problem is that some objects are older than the age of the universe at its redshift $z$. In previous literatures, many cosmological models have been tested by the old quasar APM $08279+5255$ with age $2.1\\pm0.3$ Gyr at $z=3.91$ \\cite{Friacas,Komossa}, such as the $\\Lambda$CDM \\cite{Friacas,Alcaniz2003a}, $\\Lambda(t)$ model \\cite{Cunha}, the interacting dark energy models \\cite{Wang}, Generalized Chaplygin gas model \\cite{Alcaniz2003b,Wang2009}, holographic dark energy model \\cite{Wei}, braneworld models \\cite{Movahed,Alam,Pires} and conformal gravity model \\cite{Yang2013}. But all of these models have a serious age problem except the conformal gravity model, which can accommodate this quasar at $3\\sigma$ confidence level \\cite{Yang2013}.   In this paper, we will use the old quasar APM 08279+5255 at redshift $z=3.91$ and the 9 extremely old Global Clusters \\cite{Ma2009,WangAj2010} to investigate the cosmic age problem in the $R_h=ct$ Universe. The data of these 9 extremely old GCs listed in Table 1 is acquired from the Beijing-Arizona-Taiwan-Connecticut system with up-to-date theoretical synthesis models. The evolutionary population synthesis modeling has become a powerful tool for the age determination \\cite{Tinsley68,Searle73}. In \\cite{Ma2009,WangAj2010}, they get the ages of those GCs by using multi-color photometric CCD data and comparing them with up-to-date theoretical synthesis models. But the ages of GCs derived by different authors based on different measurements using same method are not always consistent \\cite{Ma2009}. We find that the 9 of those GCs can give stronger constraints on the age of universe than the old quasar APM $08279+5255$. Those 9 extremely old GCs have been used to test the cosmic age problem in dark energy models in previous work and many dark energy models have a serious age problem \\cite{Wang}. The $R_h=ct$ Universe is a cosmic model which is closely restricted by the cosmological principle and Weyl's postulate \\cite{Melia2007}. In the $R_h=ct$ Universe, the gravitational horizon $R_h$ is always equal to $ct$. The $R_h=ct$ Universe can fit the SNe Ia data well \\cite{Melia2012} , explain the growth of high-$z$ quasars \\cite{Melia2013}, account for the apparent absence in CMB angular correlation \\cite{Melia2014}. As we discuss above, many cosmological models can not pass the age test. But whether the $R_h=ct$ Universe suffers the cosmic age problem in still unknown.  The structure of this paper is as follows. In section 2, we introduce the $R_h=ct$ Universe. In section 3, we give the constraints on the $R_h=ct$ Universe from SNe Ia and H(z) data. Then we will test the $R_h=ct$ Universe with the 9 extremely oldest GCs and the old quasar APM $08279+5255$. The age test in other cosmological models is presented in section 4. Conclusions will be given in section 5. ",
        "conclusions": "\\label{sec:Conclusions}  In this paper, we test the cosmic age problem in several cosmological models by using nine extremely old GCs in M31 and the old quasar APM $08279+5255$. We find that the best-fit value of Hubble constant in the $R_h=ct$ Universe is $H_0 = 70.01\\pm0.40 {\\rm~km~s^{-1}~Mpc^{-1}}$ at $1\\sigma$ confidence level by using SNe Ia data. In this case, the age of local $R_h=ct$ universe $t_0 = 13.97{\\pm}0.08 \\rm Gyr$. If we fit the $R_h=ct$ Universe with the SNe Ia and H(z) data, the Hubble constant is $H_0 = 69.83\\pm0.40{\\rm~ km~s^{-1}~Mpc^{-1}}$ at the $1\\sigma$ confidence level. The age of local universe is $t_0 = 14.01{\\pm}0.08$ Gyr. From Fig. \\ref{fig:Age_Ia_Hz_Q_fig}, we find that the $R_h=ct$ Universe can accommodate the old quasar APM $08279+5255$ at more than 3$\\sigma$ confidence level. From Fig. \\ref{fig:Age_Ia_Hz_G_fig}, we find that there are five GCs (B239, B144D, B260, B383, B495) can be accommodated by the $R_h=ct$ Universe at $1\\sigma$ confidence level. But the age estimates of some GCs are controversial. For example, the metallicities of B129, B024, B297D and B050 measured by \\cite{Ma2009,WangAj2010} and \\cite{Galleti} are significantly different. So the derived ages are different. Due to the uncertainty of age determination, we can claim that the $R_h=ct$ Universe can marginally solve the cosmic age problem.  Using the same method, we also test some other cosmological models, such as $\\Lambda$CDM, interacting dark energy model, generalized Chaplygin gas model and holographic dark energy model. In Sec.(\\ref{sec:OtherModels}), we show that these models can not accommodate all nine old GCs in M31. Meanwhile, for the old quasar APM $08279+5255$ at $z=3.91$, the $R_h=ct$ model can accommodate it at more than $3\\sigma$ confidence level. But these models can not accommodate it. The generalized Chaplygin gas model is in tension (over $2\\sigma$ confidence level) with the age of APM $08279+5255$. So the $R_h=ct$ Universe can marginally solve the cosmic age problem, especially at high redshift."
    },
    "1402/1402.6743_arXiv.txt": {
        "abstract": "We report the discovery of 24 Lyman-break candidates at $7\\lesssim z \\lesssim 10.5 $, in the Hubble Frontier Fields (HFF) imaging data of Abell 2744 ($z=0.308$), plus \\emph{Spiter}/IRAC data and archival ACS data. The sample includes a triple image system with a photometric redshift of $z\\simeq 7.4$. This high redshift is geometrically confirmed by our lens model corresponding to deflection angles that are 12\\% larger than the lower-redshift systems used to calibrate the lens model at $z= 2.019$.  The majority of our high-redshift candidates are not expected to be multiply lensed given their locations in the image plane and the brightness of foreground galaxies, but are magnified by factors of $\\sim 1.3 - 15$, so that we are seeing further down the luminosity function than comparable deep field imaging. It is apparent that the redshift distribution of these sources does not smoothly extend over the full redshift range accessible at $z<12$, but appears to break above $z=9$. Nine candidates are clustered within a small region of $20\\arcsec$ across, representing a potentially unprecedented concentration. Given the poor statistics, however, we must await similar constraints from the additional HFF clusters to properly examine this trend. The physical properties of our candidates are examined using the range of lens models developed for the HFF program by various groups including our own, for a better estimate of underlying systematics.  Our spectral-energy-distribution fits for the brightest objects suggest stellar masses of $\\simeq 10^{9}$~\\Msun, star-formation rates of $\\simeq 4$~\\Msun~yr$^{-1}$, and a typical formation redshift of $z\\lesssim 19$. ",
        "introduction": "Our understanding of the first few billion years of cosmic time has increased significantly in recent years, thanks to the {\\em Hubble Space Telescope's} (\\HST) Wide-Field Camera~3/Infrared Channel \\citep[WFC3/IR,][]{wfc3} as well as the {\\em Spitzer Space Telescope's} Infrared Array Camera \\citep[IRAC,][]{irac}. Until recently, the Hubble Ultra Deep Field \\citep{beckwith,garth} has provided our deepest view of the Universe, revealing a considerable number of galaxy candidates at $z > 7$, including one candidate at $z\\simeq 10$ \\citep{bouwens, bouwens1, ellis, oesch, garth}.  The cosmic epoch of $z\\simeq 10$ is important to study as it marks the dawn of galaxy formation and the beginning of reionization of the intergalactic medium.  However, galaxies at that redshift are extremely faint, making it difficult to discover and study the abundant population of galaxies below $L^{*}$, the knee of the luminosity function.  Fortunately, the magnification boost afforded by gravitational lensing combined with \\HST's exquisite imaging capabilities in the near-infrared (NIR), provides an avenue for both discovering and characterizing the intrinsic properties of galaxies around $z\\simeq 12$, when the Universe was less than half a billion years old.  The Cluster Lensing And Supernova survey with Hubble \\citep[CLASH,][]{postman} carried out \\HST{} imaging of 25 galaxy clusters in 16 broad bands between $0.2-1.7$ $\\mu$m to a depth of AB magnitude $\\sim 27$ with a total of 20 orbits per cluster.  The CLASH program has led to many interesting discoveries of magnified, intrinsically faint galaxies. Several hundred dropout galaxies have been uncovered in the range $z\\simeq 6-8$, with a few notable examples at higher redshifts of $z\\simeq 9-11$ \\citep[see][]{zheng, bouwens2, coe, bradley}, helping to motivate dedicated deeper lensing surveys.  The Hubble Frontier Fields (HFF) is a new initiative now being carried out to observe the distant Universe to an unprecedented depth, combining the power of deep \\HST\\ imaging and gravitational lensing. In \\HST's Cycles 21 and 22, 560 orbits of Director's Discretionary Time have been allocated to observe four clusters. The observations are carried out with four WFC3/IR filters (F160W, F140W, F125W, F105W) and three ACS filters \\cite[Advanced Camera for Surveys,][F814W, F606W, F435W]{acs}.  It is anticipated that 280 orbits will be allocated in Cycle 23 to observe two additional clusters. In addition, deep \\emph{Spitzer} and \\emph{Chandra} observations are planned for the six HFF fields.  These coordinated observations will enable us to probe the star formation rate density at $z\\gtrsim 9$, study the faint end of the galaxy population at $z \\simeq 3-8 $, and map the dark matter in these clusters in unprecedented detail via many multiple images of background sources (Hubble Deep Fields Initiative 2012 Science Working Group Report).\\footnote{http://www.stsci.edu/hst/campaigns/frontier-fields/HDFI\\_SWGReport2012.pdf}  We report the discovery of 24 candidate Lyman-break galaxies (LBGs) at $z\\gtrsim 7$ in the field of \\cl, based on the HFF observations and archival data. The faintest sources detected are around AB magnitude 29. These objects, listed in Tables~\\ref{tbl-z9}-\\ref{tbl-z7}, have ``secure'' photometric redshifts greater than 7 and a negligible probability ($<1\\%$) of being at lower redshift.    We adopt a concordance cosmology with $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $h=H_0/100\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}=0.7$, and the AB magnitude system throughout. ",
        "conclusions": "We find 24 LBG candidates at $7 \\lesssim z \\lesssim 10.5$ in the HFF imaging of \\cl, reaching an intrinsic magnitude of $\\sim 32$. One source at $z\\simeq 7.4$ is lensed into three images. Significant clustering is observed on the intrinsic scale of $10-100$~kpc. Thanks to gravitational lensing, we are able to carry out Spitzer/IRAC photometry for 16 of the sources. SED fitting to the brightest candidates suggests stellar masses of $\\simeq 10^9$~\\Msun, star-formation rates of $\\simeq 4$~\\Msun\\ per year, and a typical formation redshift of $z \\simeq 19$.  The redshift distribution of our sample is not a smoothly declining function towards higher redshift. In particular, our redshift distribution does not extend smoothly beyond $z\\simeq 9$, clustering notwithstanding, Considering the effect of cosmic variance, the number density in our sample is consistent with that derived from other studies, {\\it e.g.}, \\cite{bouwens4}. Given the level of clustering that we see in \\cl, it will be  important to average over more HFF, and to perform a luminosity function analysis so that the redshift dependence can be better related to galaxy mass.  The work presented in this paper is based on observations made with the NASA/ESA {\\it Hubble Space Telescope}, and has been supported by award AR-13079 from the Space Telescope Science Institute (STScI), which is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS~5-26555.  It is also based on data obtained with the {\\it Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.  This work utilizes gravitational lensing models produced by P.I.s Bradac, Ebeling, Merten \\& Zitrin, Sharon, and Williams, funded as part of the \\HST\\ Frontier Fields program conducted by STScI.  These models were calibrated using arcs identified in archival \\HST{} imaging by \\cite{merten}, spectroscopic redshifts of arcs obtained using the VLT/FORS2 spectrograph (J. Richard et~al. 2014, in prep.), and VLT and Subaru/Suprimecam imaging of the Abell~2744 field \\citep{cypriano, okabe08, okabe10, okabe10b}.  XS acknowledges support from FP7-SPACE-2012-ASTRODEEP-312725, and NSFC grants 11103017 and 11233002.  Support for AZ is provided by NASA through Hubble Fellowship grant HST-HF-51334.01-A awarded by STScI. FEB acknowledges support from Basal-CATA PFB-06/2007, CONICYT-Chile grants (FONDECYT 1141218, ALMA-CONICYT 31100004, Gemini-CONICYT 32120003, Anillo ACT1101), and the Millennium Institute of Astrophysics (Project IC120009). JMD acknowledges support from the Spanish consolider project CAD2010-00064 and AYA2012-39475-C02-01. We thank M. Meneghitti for helpful comments.  \\hskip 1.5in"
    },
    "1402/1402.3166_arXiv.txt": {
        "abstract": "We present optical photometric monitoring of KT Eri (Nova Eridani 2009), a He/N very fast nova which oubursted in November 2009. Our observations include $BVR_cI_c$ brightness estimations as well as monitoring of the rapid brightness variations in V band. The characteristic times of these rapid changes are studied and compared with the observed in other novae. ",
        "introduction": "KT Eri was discovered by \\citet{2009CBET.2050....1I} in November 2009 at high galactic latitude. Due to this unusual position it is the first classical nova observed in the constellation Eridanus. \\citet{2009ATel.2327....1R} estimated that the maximum brightness $V\\sim5.\\!\\!^\\mathrm{m}6$ occurred on November 15 which shows that the nova is very fast with $t_\\mathrm{2}\\sim8$\\,days and $t_\\mathrm{3}\\sim15$\\,days. \\citet{2009CBET.2053....3M} classified KT Eri as He/N nova. The pre-outburst light curve shows large variations in the optical and an average brightness $\\sim 15^\\mathrm{m}$ \\citep{2009ATel.2331....1D}. Because of the relatively low amplitude ($\\sim 10^\\mathrm{m}$), unusual for very fast classical novae, it was suggested \\citep{2010ApJ...724..480H} that KT Eri is in fact recurrent nova, but no evidence of other outbursts are yet available. Studying the historical light curve of the nova, \\citet{2012A&A...537A..34J} found a possible orbital period of 737 days and suggested that the secondary component could be a RGB star.  \\hspace{1em} Quasi-periodic Oscillations (QPOs) are rapid brightness variations often observed in cataclysmic variables. Their typical periods are from 3s to $\\sim$1000s usually unstable and with low amplitudes. QPOs with short periods are more profoundly studied, although all QPOs nature is uncertain (see \\citeauthor{2004PASP..116..115W} \\citeyear{2004PASP..116..115W}, \\citeauthor{2008AIPC.1054..101W} \\citeyear{2008AIPC.1054..101W}).  \\hspace{1em} \\citet{2010ATel.2392....1B} were the first to report on possible rapid brightness variations in optical observations of  KT Eri obtained from day 17 to day 65 after the maximum. Unfortunately, the frequency of their data sampling was too small for periodogram analysis. Very short $P\\sim35.09$ sec QPO were observed by \\citet{2010ATel.2423....1B} in X-ray between 67 and 79 days after the maximum. Long-term variability of KT~Eri with a period of 56.7 days was reported by \\citet{2011ASPC..451..271H}, interpreted by them as rotation of the hot spot on the accretion disk. ",
        "conclusions": "The results of our time-resolved photometry show that the KT Eri post-outburst optical QPOs appear and disappear in several day intervals. Taking into account the accuracy of our observations we cannot be sure that the QPOs really are absent. It is possible that we are not able to detect the ones with lower amplitudes.  \\hspace{1em} There are several novae which show a variety of QPO periods: IX Vel $\\sim$500 sec, V533 Her $\\sim$1400 sec, BT Mon $\\sim$1800 sec, GK Per $\\sim$5000 sec, V842 Cen $\\sim$750-1300 sec (Warner 2004), V2468 Cyg $\\sim$1260-3000 sec \\citep{2012MmSAI..83..767C}. Most of these periods, excepting GK Per QPOs period are shorter in comparison to ones of KT Eri.  \\hspace{1em}  It is interesting to note, that about two years after the outburst of the recurrent nova RS Oph in 2006,  \\citet{2009ASPC..404...72V} found optical QPOs with a period $\\sim$3000 sec. Very close to the observed in KT Eri on days 133 and 142.  Furthermore,  \\citet{2011ApJ...727..124O} detected X-ray QPOs with a period $\\sim$35 sec during the supersoft phase of RS Oph, similar to the reported by Beardmoreday et al. (2010) for KT Eri.  \\hspace{1em} The resemblance of the KT Eri and RS Oph quasi periodic oscillations may be an additional argument in favor of the suggestion that KT Eri is in fact a recurrent nova. But the large variety of the periods and the low stability of the QPOs observed in novae make this argument insignificant."
    },
    "1402/1402.5095_arXiv.txt": {
        "abstract": "In this paper we use near-infrared (NIR) spectral observations of Type~Ia supernovae (SNe~Ia) to study the uncertainties inherent to  NIR $K$ corrections. To do so, 75 previously published NIR spectra of 33 SNe~Ia are employed to determine $K$-correction uncertainties in the $YJHK_s$ passbands as a function of temporal phase and redshift. The resultant $K$ corrections are then fed into an interpolation algorithm that provides mean $K$~corrections as a function of temporal phase and robust estimates of the associated errors.  These uncertainties are both statistical and intrinsic --- i.e., due to the diversity of spectral features from object to object --- and must be included in the overall error budget of cosmological parameters constrained through the use of NIR observations of SNe~Ia.  Intrinsic variations are likely the dominant source of error for all four passbands at maximum light.  Given the present data, the total $Y$-band $K$-correction uncertainties at maximum are smallest, amounting to $\\pm 0.04$~mag at a redshift of $z = 0.08$.  The $J$-band $K$-term errors are also reasonably small ($\\pm 0.06$~mag), but intrinsic variations of spectral features and noise introduced by telluric corrections  in the $H$-band currently limit the total $K$-correction errors at maximum to $\\pm 0.10$~mag at $z = 0.08$.  Finally, uncertainties in the $K_s$-band $K$~terms at maximum amount to $\\pm 0.07$~mag at this same redshift. These results are largely constrained by the small number of published NIR spectra of SNe~Ia, which do not yet allow spectral templates to be constructed as a function of the light curve decline rate. ",
        "introduction": "Type Ia supernovae (hereafter SNe~Ia) are standardizable distance indicators at optical wavelengths that provide critical constraints on cosmological parameters \\citep[see][and references therein]{goobar11}. A number of groups have worked diligently to gather optical photometry of homogenous samples of low-, intermediate- and high-$z$  SNe~Ia that, when combined, amount to  well over 1000 objects. As the sample size has increased, systematic effects have come to dominate the final uncertainty in the measured value of the equation-of-state parameter of the Universe, $w$ \\citep[e.g., see][]{wood07,astier06,freedman09,kessler09,folatelli10,conley11,suzuki12}.  A major systematic  that  plagues SN~Ia cosmology is our inability to accurately  estimate host galaxy dust extinction due to uncertainties in the  reddening law and, in particular, variations in the value of the total-to-selective extinction, $R_V$, from object to object \\citep[see][and references therein]{phillips12}. This is further exacerbated by any systematic error in the relative zero points between the nearby and distance SNe~Ia samples, thereby making the latter artificially redder (or bluer) than the former. A sensible way around these problems is to observe in rest-frame near-infrared (NIR) bands instead of rest-frame optical bands, because the effects of dust extinction are minimized and essentially independent of the adopted reddening law \\citep{krisciunas00}.  Additional motivation is provided by empirical evidence indicating that the luminosities of SNe~Ia show little or no dependence on decline rate\\footnote[8]{The decline rate of a SN~Ia is traditionally defined as the change in its $B$-band magnitude from the time of maximum brightness to 15 days later, and is denoted as  $\\Delta$$m_{15}(B)$.   The decline rate is known to correlate with the peak absolute luminosity in such a way that more luminous objects exhibit smaller  $\\Delta$$m_{15}(B)$  values \\citep{phillips93}.} in the NIR \\citep{meikle00,kr204,wood08,kriciunas09,mandel09,folatelli10,kattner12}. This translates to a reduced intrinsic dispersion in the NIR Hubble diagram, and at the same time evolutionary effects as a function of redshift on the progenitor populations are potentially minimized. These  factors have provided significant impetus for future SN~Ia cosmology studies to construct {\\it homogenous samples} of low-  and high-$z$  SN~Ia observed in the rest-frame NIR \\citep[e.g., see][]{green12}.  Reaping the benefits afforded by observing SNe~Ia in the NIR requires the development of tools to obtain rest-frame luminosities via $K$~corrections \\citep{oke68}. The $K$~correction is defined as the difference in brightness between an object observed in its rest-frame with a given passband compared to its measured brightness with the same passband when observed at redshift $z$. Specifically, for a given passband, $i$, the $K$~correction is defined as: \\begin{equation} m_{i}=M_{i} + \\mu + K_{i}. \\label{K_corr1} \\end{equation} Here $m_{i}$ is  the SN~Ia apparent magnitude observed on Earth, and $M_{i}$ is its absolute magnitude. The magnitude difference encapsulated in the $K$ term is explained by the shifting and stretching of an object's spectral energy distribution  (SED), which is inherent to cosmological expansion.  The first calculations of SN~Ia $K$~corrections at optical wavelengths were published by \\citet{leibundgut90} and \\citet{ham93}.  In the latter paper, optical observations of three nearby SNe~Ia were used to construct a sequence of $B$- and $V$-band $K$~corrections extending to a redshift of $z = 0.5$. The authors found that the temporal variation of the computed values largely mimicked the $(B-V)$ color evolution, implying that the $K$~correction is, to first order, driven by the color of the SN.  Shortly thereafter,  \\citet{kim96} presented a method to compute cross-band $K$~corrections which is  particularly well-suited for SNe~Ia located at $z > 0.2$.   \\citet{nugent02} expanded upon these efforts by developing a set of SN~Ia  optical spectral templates that enabled $K$~corrections to be computed as a function of  temporal phase for any given optical bandpass. Later, \\citet{hsi07} constructed improved optical spectral templates based on a greatly expanded sample of nearby SN~Ia.  They found that besides color, spectral diversity also has a significant impact upon the magnitude of the $K$~term. The \\citet{hsi07} template is now routinely used to $K$~correct optical photometry of SNe~Ia, and the statistical uncertainty associated with these corrections is reasonably well understood.  In comparison, our knowledge of NIR $K$~corrections is still relatively crude. \\citet{krisciunas04} presented the first calculations of the temporal evolution of the $K$~term in the $JHK$ passbands based on 11 NIR spectra of SN~1999ee  \\citep{hamuy02}. Their results indicated that NIR $K$~corrections are non-negligible even at relatively small redshifts. These initial findings emphasized  the importance of accurately characterizing NIR $K$~corrections {\\it and} their uncertainties as a function of redshift and temporal phase. Five years later, the publication of 41 spectra by \\citet{marion09} revolutionized the study of the NIR spectral characteristics of SNe~Ia, and \\citet{hsiao09} used these data along with the other published spectra available at that time to extend his spectral template to include NIR wavelengths.  In this paper, we refer to this template as the ``Hsiao revised template''.  In the present paper, we expand on this work by using the existing library of published NIR SNe~Ia spectra to determine the uncertainties inherent to using the Hsiao revised spectral template to calculate NIR $K$~corrections, particularly those due to intrinsic variations in spectral features. To do this, we first color match\\footnote[9]{Often people tend to use the nomenclature warp or mangle, however, we feel that it is more appropriate and accurate to adopted the term color match.} each observed spectrum to the template, and then calculate  $YJHK_s$-band $K$~corrections. An interpolation algorithm based on Gaussian Processes combined with Markov-Chain Monte-Carlo methodology is then used to produce mean $K$~corrections as a function of temporal phase and redshift, along with estimates of both statistical and intrinsic sources of uncertainties. These errors, in turn, will provide important input to future studies utilizing the NIR light curves of SNe~Ia to estimate cosmological parameters.  The organization of this article is as follows. In Section~\\ref{kcorrection}  the concept of the $K$~correction is briefly reviewed; Section~\\ref{se:data} introduces the data used in our calculations, along with the adopted passbands and methods used to interpolate our computed $K$~corrections; Section~\\ref{se:results} contains the results; and finally, Section~\\ref{conclusions} presents our conclusions. ",
        "conclusions": "\\label{conclusions}  Using a library of publicly available SNe~Ia NIR spectra, we have computed a set of  $K$~corrections in the CSP $YJHK_s$ filters at redshifts of  $z =0.03$, 0.05, and 0.08. The individual spectra were first color-matched to the Hsiao revised spectral template before calculating the $K$~corrections. A combined Gaussian Process and MCMC method was then employed to derive an empirically-based model of the $K$~terms as a function of temporal phase and redshift.  This procedure returns uncertainties in the $K$-correction model due to both statistical noise and intrinsic diversity in the features of the input spectra.  \\citet{krisciunas04} published a table of  NIR $K$ corrections computed using spectra of SN~1999ee,  and a set of filters functions similar to those adopted in this study. A comparison between the results of \\citeauthor{krisciunas04}  to those obtained with our expanded set of spectra and more advanced color-matching technique provides good agreement in the $J$ and $H$ bands at $z=0.03$, while for the $K_s$ band we find differences of up to $\\pm$0.08 mag. The discrepancy in the $K_s$ is not surprising owing the the dearth of spectra included in the work of \\citeauthor{krisciunas04}, as well as the difficulties highlighted in this study concerning the difficulties inherent to computing robust $K_s$-band $K$ corrections.   Our emphasis in this paper has not been to provide ``lookup\" tables of NIR $K$~corrections, but rather to derive  accurate estimates of the uncertainties inherent in the use of  spectral templates (like the Hsiao revised template) constructed from the NIR sample of SNe~Ia spectra available today.  We find that the $Y$ and $J$ bands currently afford the greatest precision in $K$-correction calculations due to the general weakness of the spectral features and the minimal effect of telluric corrections.  The $Y$ band is particularly noteworthy, with the uncertainty due to spectral feature variations increasing from only $\\pm0.02$ to $\\pm$0.04~mag from $z = 0.03$--0.05. The $H$ band is currently more problematic due to noise introduced by telluric corrections and the appearance just after maximum light of strong Fe-peak emission features. Diversity in the strength of this emission most likely explains the strong increase we find in the intrinsic component of the $K$-correction error from $\\pm0.02$ to $\\pm$0.10~mag over the redshift range $z = 0.03$--0.05. This result  is a departure from pervious theoretical \\citep{kasen06} and observational \\citep{wood08} findings, which have identified the $H$ band as having potentially the least dispersion among the NIR passbands. Finally, the uncertainties in the $K$~terms for the $K_s$ filter due to intrinsic spectral diversity are similar to those found for the $J$~band, but the currently-available library of spectra covering this wavelength region suffer from both low signal-to-noise and telluric absorption.  Further progress in reducing the uncertainties in the $K$~corrections for NIR filter bandpasses demands many more spectral observations over a wider range of redshifts. Figure~\\ref{hist} shows a histogram of the decline rates of the SNe~Ia in our sample with temporal phases within $\\pm 3$~days of  \\bmax.  For reference, the same figure includes a plot of the decline rates of the spectra employed by \\citet{hsi07} to create their optical spectral template.  The discrepancy between these two histograms dramatically illustrates the limitations of the currently-available library of NIR spectra. To address this problem, we have embarked on a four-year, six-months-a-year SN~Ia followup program that builds upon the legacy of the CSP. This project, designated CSP-II,  is designed to obtain optical and NIR  light-curves of $\\sim$100 SNe~Ia in the smooth Hubble flow ($0.03 < z < 0.08$).  Complementary to these observations, frequent NIR spectroscopy is being carried out of nearby SNe~Ia, mainly with the Folded-port InfraRed Echellette (FIRE) spectrograph mounted on the 6.5-m Magellan Baade telescope. Particular emphasis is being placed on obtaining spectral sequences for a sub-sample of the SNe that span a range in decline rate and phase, which are essential to determining correlated errors in the $K$~corrections. We are confident that such observations will not only allow us to more precisely characterize NIR $K$-correction uncertainties, but also to decrease these errors to the low levels now obtained at optical wavelengths.  In order to ascertain the contribution of NIR $K$-correction errors to the final error budget of cosmological parameters one must determine whether  they are systematic in temporal phase, and therefore propagate to the peak magnitude of a SN~Ia. On the other hand, if the $K$-correction errors are random in phase then the final uncertainty in the peak magnitude will be reduced by template fitting. To determine which of these is the case is, however, beyond the scope of this paper due to the lack of a statistical sample of SNe~Ia with good spectral time series. However, in the near future we will be able to tackle this issue with the high-fidelity spectral series currently being obtained."
    },
    "1402/1402.2483_arXiv.txt": {
        "abstract": "{ L1642 is one of the two high galactic latitude ($|b| > 30^{\\circ}$) clouds confirmed to have active star formation. } { We examine the properties of this cloud, especially the large-scale structure, dust properties, and compact sources in different stages of star formation. } { We present high-resolution far-infrared and submillimetre observations with the \\emph{Herschel} and AKARI satellites and millimetre observations with the AzTEC/ASTE telescope, which we combined with archive data from near- and mid-infrared (2MASS, WISE) to millimetre wavelength observations (Planck). } { The \\emph{Herschel} observations, combined with other data, show a sequence of objects from a cold clump to young stellar objects at different evolutionary stages. Source B-3 (2MASS J04351455-1414468) appears to be a YSO forming inside the L1642 cloud, instead of a foreground brown dwarf, as previously classified. \\emph{Herschel} data reveal striation in the diffuse dust emission around the cloud L1642. The western region shows striation towards the NE and has a steeper column density gradient on its southern side. The densest central region has a bow-shock like structure showing compression from the west and has a filamentary tail extending towards the east. The differences suggest that these may be spatially distinct structures, aligned only in projection. We derive values of the dust emission cross-section per H nucleon of $\\sigma_e(250 \\mu {\\rm m}) = $ 0.5--1.5 $\\times10^{-25} {\\rm cm}^2/{\\rm H}$ for different regions of the cloud. Modified black-body fits to the spectral energy distribution of \\emph{Herschel} and Planck data give emissivity spectral index $\\beta$ values 1.8--2.0 for the different regions. The compact sources have lower $\\beta$ values and show an anticorrelation between $T$ and $\\beta$. } { Markov chain Monte Carlo calculations demonstrate the strong anticorrelation between $\\beta$ and $T$ errors and the importance of millimetre wavelength Planck data in constraining the estimates. L1642 reveals a more complex structure and sequence of star formation than previously known. } ",
        "introduction": "\\label{sect:intro}  High galactic latitude ($|b| > 30^{\\circ}$) molecular clouds have low background and foreground emission, and are usually quite nearby. They are mostly diffuse or translucent clouds with a low density and typically no signs of star-formation, see~\\citet{McGehee2008} for a review. These features make them ideal for studying the interstellar medium and interstellar radiation field, as well as low-mass objects in selected regions. LDN 1642\\footnote{As a side note, a memory aid for the number: in the year 1642 Galileo Galilei died and Sir Isaac Newton was born.}~\\citep{Lynds1962} (often also called L1642) is one of the two clouds at high galactic latitudes confirmed to have active star-formation, the other one is MBM 12~\\citep{Luhman2001}. Studying this cloud might give constraints of the mechanisms that affect cloud evolution and trigger low-mass star formation. In this article, we use the abbreviation L1642. The cloud is also called MBM 20~\\citep{Magnani1985} and G210.90-36.55~\\citep{Juvela2012}.  L1642 is one of the Orion outlying clouds, see~\\citet{Alcala2008} for a review of the area and earlier studies. It forms a blob in the head of a HI cloud with cometary structure, the tail of which extends over 5$^{\\circ}$ in the north-east direction in equatorial coordinates, directly towards the Galactic plane. The cloud is projected on the edge of the Orion-Eridanus Bubble~\\citep{Brown1995}. On the sky, L1642 is $\\sim10^{\\circ}$ from the filamentary reflection nebula IC 2118, or Witch Head nebula, which also contains the molecular clouds MBM 21 and 22~\\citep{Kun2001,Kun2004,Alcala2008}. The different clouds are marked in Fig.~\\ref{fig:Planck_Ak}.  The Galactic coordinates of L1642 are $l = 210.9^{\\circ}$ and $b = -36.55^{\\circ}$, and the equatorial coordinates are $\\alpha_{2000} = 4^{\\rm h}35^{\\rm m}$ and $\\delta_{2000} = -14^{\\circ}15'$ ($\\sim 68.75^{\\circ}$ and $-14.25^{\\circ}$, respectively). \\citet{Hearty2000} determined the distance of L1642 to be 112--160 pc. The X-ray observations of~\\citet{Kuntz1997} suggested that the cloud is close to the edge of the Local Bubble, and not necessarily inside the Orion-Eridanus Bubble, indicating a distance of approximately 140 pc according to \\citet{Sfeir1999}. See~\\citet{Welsh2009} for a description of the Local Bubble. We adopted the distance of 140 pc similarly to~\\citet{Russeil2003} and~\\citet{Lehtinen2004}.  L1642 has been studied using several molecules, see~\\citet{Liljestrom1991} for a review of the early studies. \\citet{Liljestrom1991} studied L1642 in CO, HCO$^+$ and NH$_3$ and concluded that the cloud is in virial equilibrium and older than $10^6$ years. More recently, \\citet{Russeil2003} studied the morphology and kinematics of L1642 using CO observations obtained with the SEST radio telescope. Their results showed that the cloud consists of a main structure at radial velocity 0.2 km s$^{-1}$ with higher velocity components forming an incomplete ring around it, suggesting an expanding shell. Based on $^{13}$CO data, the peak column density is $N$(H$_2$) $\\sim6\\times10^{21}$ cm$^{-2}$ and the cloud mass $\\sim$59 $M_{\\odot}$.  L1642 has also been widely studied using other methods such as dust emission in mid- and far-infrared~\\citep{Laureijs1987,Reach1998,Verter1998,Verter2000,Lehtinen2004,Lehtinen2007} and optical (360--600 nm) surface brightness~\\citep{Laureijs1987,Mattila2012}. \\citet{Mattila2007} studied scattered H{$\\alpha$} light in L1642. The shadowing provided by L1642 has also been used in studies of Galactic diffuse X-ray radiation~\\citep[e.g.,][]{Galeazzi2007,Gupta2009}.  The small distance, high galactic latitude with little foreground contamination, and the modest column density make L1642 a good target for studying low-mass star formation and the effect of potential external triggering. \\citet{Sandell1987} identified two IRAS sources, IRAS 04327-1419 = L1642-1 (V* EW Eri, HBC 413) and IRAS 04325-1419 = L1642-2 (HBC 410), found within L1642, to be faint nebulous binary stars with very active secondaries. The primary of L1642-1 was spectroscopically classified to be a K7IV T Tauri star. \\citet{Liljestrom1989} found a weak, bipolar outflow around the binary L1642-2. \\citet{Reipurth1990} discovered a Herbig-Haro object, HH123, also originating from L1642-2. They also concluded that the primary object of L1642-2 is a low-luminosity M0 H$\\alpha$-emission star, and the secondary component is an H$\\alpha$-emission star as well. \\citet{Correia2006} observed L1642-1 in search of high-order multiplicity, but did not find a third component. \\citet{Lehtinen2004} concluded that none of the other four IRAS sources projected within L1642 are likely to be young stellar objects (YSOs) inside the cloud.  The 2MASS point source catalogue object 2MASS J04351455-1414468 is situated towards the densest part of L1642 on the sky. \\citet{Cruz2003} classified it as a young ($\\sim$10 Myr) object. They estimated the distance to be probably within 30 pc, meaning that the object would be clearly outside the L1642 star-forming region. Later studies have made the assumption or conclusion that the object is a brown dwarf located in front of the L1642 cloud, at a distance of 14--30 pc~\\citep{Faherty2009,Antonova2013}. In this article, we re-evaluate the classification and distance estimate of this object.  The early low-resolution studies treated the L1642 cloud as a single entity. \\citet{Lehtinen2004} presented ISOPHOT far-infrared (FIR) data of L1642 and studied the cloud structure in separate components (named A1, A2, B, and C). Neither of the IRAS point sources nor new YSO candidates were seen in their 200 $\\mu$m maps. \\citet{Lehtinen2007} compared far-IR data with visual extinction and studied the dust grain emissivity properties in L1642. They found that FIR dust emissivity increases by a factor of two between regions with colour temperatures 19 K and 14 K. This might be caused by grain growth in the dense and cold areas.  L1642 was recently observed as part of the \\emph{Herschel} open time key programme Galactic Cold Cores~\\citep{Juvela2010} with unprecedented spatial resolution and sensitivity at 100, 160, 250, 350 and 500 $\\mu$m. \\citet{Juvela2012} presented these observations and studied the properties of the cloud on large-scale images, without separating the distinct objects. In Rivera-Ingraham et al. (in prep) we will be carrying out a detailed \\emph{Herschel}-based compact source and environmental study of all high-latitude fields (including L1642) from our programme. Furthermore, we have carried out 1.1 mm dust-continuum observations toward L1642 with the AzTEC/ASTE instrument. In this paper, we use these new data, along with other, already published data, to complete the spectral energy distribution of L1642 in the long-wavelength region and advance with this (i) the study of the large-scale structure of the cloud, (ii) the investigation of general dust properties, and (iii) the characterisation of point and compact sources related to the cloud.  The contents of the article are the following: we present the observations and data processing in Sect.~\\ref{sect:observations} and methods in Sect.~\\ref{sect:methods}. The results based on the observations are shown in Sect.~\\ref{sect:results}. We present a radiative transfer model of L1642 in the appendix. We discuss the results in Sect.~\\ref{sect:discussion} and present our conclusions in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have studied the properties of L1642, one of the two high galactic latitude ($|b| > 30^{\\circ}$) clouds with active star-formation. The high latitude ($-$36.4$^{\\circ}$), indicating no foreground contamination, the small distance ($\\sim$140 pc), the modest column density ($A_V < 12^{\\rm m}$), and the spatially distinct sequence of star-forming cores make L1642 one of the best targets for studying the initial stages of low-mass star formation. We presented in detail high-resolution far-infrared and submillimetre observations with \\emph{Herschel} and AKARI satellites and millimetre observations with AzTEC/ASTE telescope and combined them with archive data from NIR (2MASS) and MIR (WISE) to millimetre (Planck). We studied both larger regions (A1, A2, B, C) and compact sources (both previously known binary T Tauri stars and potential new objects) within the main cloud. Our main conclusions are the following:  \\begin{itemize}  \\item The high-resolution \\emph{Herschel} observations, combined with other data, show an evolutionary sequence from a cold clump (B-4) to young stellar objects of different spectral classes: B-1 (L1642-1, Class II), B-2 (L1642-2, Flat spectrum), and B-3 (Class III).  \\item Based on \\emph{Herschel} FIR--submm data, the point source B-3 (2MASS J04351455-1414468) appears to be a YSO forming inside the L1642 cloud, instead of a foreground brown dwarf, contrary to previous classification.  \\item \\emph{Herschel} data reveal striation in the diffuse dust around the cloud L1642, especially in the northern part. Region A shows striation extending towards the NE and has a correspondingly steeper column density gradient on its southern side. Region B has the appearance of being compressed from the west and has a filamentary tail extending eastward. The differences suggest that these may be spatially distinct structures, aligned only in projection.  \\item Both \\emph{Herschel} FIR--submm and AzTEC/ASTE mm data show an elongated structure north-east of the binary star B-2 (L1642-2), which was not visible in earlier observations. The structure is aligned with the general striation pattern, but it has a significantly higher column density.  \\item We derived values of the dust emission cross-section per H nucleon of $\\sigma_e(250) = $ 0.56, 0.47, 1.45, and 1.04 $\\times10^{-25} {\\rm cm}^2/{\\rm H}$ for the different regions.  \\item Modified black-body fits to the SED of \\emph{Herschel} and Planck data give emissivity spectral index $\\beta$ values 1.8--2.0 for the different regions of L1642.  \\item Markov chain Monte Carlo fits of the SEDs show a strong anticorrelation between $\\beta$ and $T$ errors and the importance of Planck data in constraining $\\beta$ and $T$ estimates.  \\item The compact sources show a $T$-$\\beta$ anticorrelation, but on larger scale the data show no clear correlation or anticorrelation.  \\item Radiative transfer modelling suggests that low beta values towards the sources might be caused to a large extent by the LOS temperature variations. On the other hand, the optical depth estimates are expected to be quite accurate (to within $\\sim$ 10 \\%) away from the sources.  \\end{itemize}"
    },
    "1402/1402.0509_arXiv.txt": {
        "abstract": "\\centering Observations of circumbinary planets orbiting very close to the central stars have shown that planet formation may occur in a very hostile environment, where the gravitational pull from the binary should be very strong on the primordial protoplanetary disk. Elevated impact velocities and orbit crossings from eccentricity oscillations are the primary contributors towards high energy, potentially destructive collisions that inhibit the growth of aspiring planets. In this work, we conduct high resolution, inter-particle gravity enabled {\\emph N}-body simulations to investigate the feasibility of planetesimal growth in the Kepler-34 system. We improve upon previous work by including planetesimal disk self-gravity and an extensive collision model to accurately handle inter-planetesimal interactions. We find that super-catastrophic erosion events are the dominant mechanism up to and including the orbital radius of Kepler-34(AB)b, making {\\it in-situ} growth unlikely. It is more plausible that Kepler-34(AB)b migrated from a region beyond 1.5 AU. Based on the conclusions that we have made for Kepler-34 it seems likely that all of the currently known circumbinary planets have also migrated significantly from their formation location with the possible exception of Kepler-47(AB)c. ",
        "introduction": "\\begin{figure*}[tbph] \\centering \\includegraphics[scale=0.45]{f1.eps} \\caption{Evolution of eccentricity waves at $t_0$ ({\\bf A}), $25\\,P_{AB}$ ({\\bf B}), $100\\,P_{AB}$ ({\\bf C}), $1000\\,P_{AB}$ ({\\bf D}). The instantaneous eccentricity of each planetesimal is shown with a grey value between black ($e=0$) and white ($e=0.02$ {\\bf A}, $e=0.1$ {\\bf B-D}). {\\bf E}-{\\bf H} show eccentricity evolution versus $a$ with $e_{ff}$ \\citep{paardekooper12} shown in grey.} \\label{fig:multi} \\vspace{+10pt} \\end{figure*}  Several planets have been discovered by Kepler in `p-type' orbits fully encompassing tight binaries \\citep{doyle11, welsh12, orosz12, schwamb13}. These planets orbit close to their host binaries ($a < 1.1$ AU) and are subject to significant perturbations which cause impact speeds to increase making accretion close to the binary difficult \\citep{marzari00,thebault06,scholl07}.  \\cite{paardekooper12} conducted numerical simulations of test particles within circumbinary planetesimal disks including a collision model from \\citet{stewart09} which allowed for accretion and erosion. The collision model provided an alternative mechanism for planetesimal growth: accretion of dust created in catastrophic collisions. However, the authors found that growth at the location of the Kepler-16 and Kepler-34 planets was difficult. Their investigation also revealed the contribution of short-period perturbations on the disk. This fast eccentricity forcing evolves more quickly than the gas damping timescale therefore the disk stays excited.  Thus, there remains a degree of uncertainty about whether observed circumbinary planets formed in-situ. In addition, the literature does not agree on the locations that could have supported the planets' growth within a circumbinary disk. For example, \\cite{meschiari12a} identify a narrow range of annuli just outside of 1 AU that could sustain planetesimal growth in Kepler-16 and \\cite{paardekooper12} find the accretion friendly zone is beyond 4 AU.  Most previous work omit inter-planetesimal gravity and/or a comprehensive collision model. In this paper, we present high resolution, 3D, inter-particle gravity (IPG) enabled {\\emph{N}}-body simulations of a circumbinary protoplanetary disk in order to address the question of where the circumbinary planets can form. We focus on the orbital dynamics, collisional evolution and physical growth of 100 kilometre-sized planetesimals in the Kepler-34 system. We consider a purely $N$-body case to isolate the effects of inter-planetsimal gravity, leaving the inclusion of a gas disk for future work.  To account for oblique, high speed impacts from orbit crossings, we use a state-of-the-art collision model \\citep[][Leinhardt et al.~in prep]{leinhardt12}. This collision model more accurately identifies regions where planetesimal growth can occur than previous work. We use statistical arguments to classify the feasibility of sustained growth events in the disk, addressing the primary question of whether in-situ growth of Kepler-34(AB)b is possible.  In section \\ref{sec:theory} we discuss the analytics of circumbinary planetesimals acting under perturbations from the stellar binary. Our numerical method and collision model are outlined in section \\ref{sec:method}. In section \\ref{sec:results} we present our results and section \\ref{sec:dis} discusses the broader implications. Conclusions are drawn in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:dis}  Kepler-34(AB)b orbits at 1.09 AU which falls in a regime where only 12$\\%$ of collisions lead to growth (Fig.~\\ref{fig:collspatevo}B). Although the eccentricity forcing falls off with $1/a^2$, $e_{ff}$ only drops to the mean initial eccentricity (e = 0.007) at 1.7 AU. However, it may not be essential to retrieve an identical environment to that of the single star case in order to sustain growth. By 1.5 AU mass changing collisions are split between partial erosion and partial accretion events, therefore, by 1.5 AU prolonged planetesimal growth becomes more sustainable. Thus, a likely scenario is that Kepler-34(AB)b formed beyond 1.5 AU and migrated inwards. It is worth mentioning that we retrieve this highly erosive disk that is unable to support sustained growth, despite our large-sized planetesimals.  \\begin{figure}[h!] \\centering \\includegraphics[scale=0.38,angle=0]{f5.eps} \\caption{Eccentricity distribution of planetesimals at 2,000 $P_{AB}$. Green points: $m \\geq 2m_o$, black points: $m<2m_o$. The red line shows contribution from perfect mergers.} \\label{fig:growth} \\vspace{+5pt} \\end{figure}  While the visual difference between the IPG and RIPG cases is small, the underlying effects propagate through to the collision outcomes. The lack of gravitational focusing with RIPG accounts for a higher proportion of oblique impacts that generally lead to an increase in hit-and-run events. While the majority of these collision are non-erosive, they do not contribute to the growth of the system. A more subtle point is that the inclusion of inter-planetesimal gravity results in collisions that would not otherwise occur in a scenario without gravitational focusing. This results in a number of low speed and accretion enabling collisions.  From these results we find that $e_{ff} > 0.01$ indicates a region of the planetesimal disk that is too perturbed to support planetesimal growth and $e_{ff} < 0.01$ may be calm enough for planetesimal accretion. If we apply this criteria to all known Kepler circumbinary planets, and assuming that this critical value applies to $e_f$ as well, only Kepler 47(AB)c could possibly have formed in-situ and all of the rest: Kepler 34(AB)b, 16(AB)b, 35(AB)b, 38(AB)b, 47(AB)b and 64(AB)b, must have formed further out in the planetesimal disk and migrated inwards to their current location. The effect of a reduced collision rate at large orbital radii on the formation timescale of planets is possibly explained by \\citet{alexander12} who suggests that circumbinary planets have more time to form than planets around single stars due to the longer lifetime of disks around short period binaries.    In this paper we presented results from high-resolution N-body simulations investigating the likelihood of in-situ formation of Kepler-34(AB)b. We began with a unperturbed disk using eccentricity and inclination distributions from numerical simulations of planetesimal evolution around single stars and allowed the system to evolve to quasi-steady state while suppressing collisional evolution of the planetesimals. Eccentricity forcing from the binary pumped up the planetesimal eccentricities to 10 times the initial mean value. After 1000 $P_{AB}$ a quasi-steady state was reached where planetesimals with similar orbital radii take on eccentricity extremes, such that both circularised and eccentric bodies can occupy the same space. This is shown by the speckled effect in Figure \\ref{fig:multi}{\\bf D}.  This new initial condition was highly perturbed, which meant many orbit crossing events occurred.  After the initial condition was reached we allowed the planetesimals to evolve collisionally using our new collision model EDACM based on \\citep{leinhardt12} which allows erosion, accretion, and bouncing events. The high eccentricity observed in the inner circumbinary disk translates into a dominance of super-catastrophic events seen in Figure \\ref{fig:collspatevo}. The impact velocity and parameter decrease with increasing orbital radius; the latter due to a narrower eccentricity dispersion resulting in a reduced number of orbit crossings. The number of high energy, erosive events decrease as a function of increasing semi-major axis. However, planetesimal growth events do not dominate the collision outcomes within the simulation domain ($a<1.5$ AU). We therefore show that the disk is a hostile environment even for our gravitationally strong 120 km planetesimals, suggesting even more difficulties for sustained accretion to occur in simulations that feature much smaller bodies.  Using statistical arguments from collisional data, in addition to physical growth rates, we find that Kepler-34(AB)b would struggle to grow in-situ. In addition, we suggest that from all the known Kepler circumbinary planets only Kepler-47(AB)c could have formed in-situ while the rest must have formed at larger $a$ where the protoplanetary disks were less perturbed by the binary stars and migrated inwards to their current location. We also show that inter-planetesimal gravity must be included in planet formation models in order to capture gravitational focusing effects that may be missed otherwise, such as low-velocity, growth-enabling impacts that may influence the outcome of the simulation.  Previous work which has attempted to hybridise a protoplanetary disk with a gas counterpart has suggested that the gas disk is similarly perturbed by a stellar binary causing further excitations to the planetesimal eccentricities, which, could further provoke unfavourable impact velocities \\citep{marzari08,marzari13,paardekooper08}. Our simulations, therefore, likely present a best-case scenario for planetesimal accretion."
    },
    "1402/1402.7329_arXiv.txt": {
        "abstract": "Measurements of the polarized radiation often reveal specific physical properties of emission sources, such as  strengths and orientations of magnetic fields offered by synchrotron radiation and Zeeman line emission, and the electron density distribution by free-free emission. Polarization-capable, millimeter/sub-millimeter telescopes are normally equipped with either septum polarizers or ortho-mode transducers (OMT) for the detection of polarized radiation.  While the septum polarizer is traditionally conceived to be limited to a significantly narrower bandwidth than the OMT, it does possess advantageous features for astronomical polarization measurements unparalleled by the OMT.  Challenging the conventional bandwidth limit, we design an extremely wideband circular waveguide septum polarizer, covering $42\\%$ bandwidth, from 77 GHz to 118 GHz, without any undesired resonance.  Stokes parameters constructed from the measured data in between 77 GHz and 115 GHz show that the leakage from $I$ to $Q$ and $U$ is below $\\pm 2\\%$ and the $Q-U$ mutual leakage below $\\pm 1\\%$.  Such a performance is comparable to other modern polarizers, but the bandwidth of this polarizer can be at least twice as wide.   This extremely wide-band design removes the major weakness of the septum polarizer and opens up a new window for future astronomical polarization measurements. ",
        "introduction": "Wide-band polarization measurements in millimeter and sub-millimeter Astronomy have always been a challenge.   Examples requiring wide-bandwidth polarimetry measurements include the continuum cosmic microwave background (CMB) polarization observation \\citep{kov02, kog03, bar05, mon06, sie07, fri09, chi10, bis11}, the synchrotron polarization observations of compact sources \\citep{dow98, ait00, cul11}, the polarization observation of thermal dust emission \\citep{laz03, bet07, hoa11} and the observation of Zeeman effects via molecular lines, such as CN and SO \\citep{bel89, cru99, shi00}. Except for some particular features of Zeeman effects, most of the above focus on linear polarization measurements.  For the continuum observation, it is desirable the available frequency bandwidth to be as wide as possible so that, on one hand, the signal-to-noise ratio can increase, and on the other hand the spectral index may be determined. For the line observation, the emission line can be red-shifted to any unknown frequency when emitted from, or absorbed at, a distant universe, and hence wide frequency coverage has a unique advantages.  While the current trend of CMB polarization experiments has shifted to the multi-pixels, silicon wafer-based, incoherent detector approach, such as the microwave kinetic inductance detector (MKID) \\citep{day03, mal10}, the conventional coherent detector approach still has its own merit, for its capability to control, detect and calibrate out the systematics.    In this regard, heterodyne polarimeters with wide bandwidths are highly desirable.   For synchrotron emitting compact sources, the interferometry array remains to be the means to reach high sensitivity, and it must adopt coherent detectors for polarization measurements.  Similar considerations apply to line emission and dust emission from compact molecular cloud cores.   All these make the conventional waveguide polarimeter device a time-honored instrument that will always be needed in future forefront millimeter/sub-millimeter telescopes.  Polarization measurements require separating the incoming radiation into two orthogonal components for the determination of Stokes parameters. Traditionally two competing devices are available for the separation of polarization, the septum polarizer and the ortho-mode transducer (OMT). An ideal septum polarizer can convert an input linear polarization wave into two circular polarization waves of equal power at two output ports. Interestingly, in a specific arrangement when the electric field of the input wave is perpendicular or parallel to the symmetric axis of the septum polarizer, the two output electric fields will be either in phase or 180 degrees out of phase, with the latter being delayed relative to the former by 90 degrees over a finite frequency interval.  It then follows that when the input is a circularly polarized wave, it will exit entirely through only one output port and the other port has a null output.  This novel feature of the septum polarizer makes it distinct from the other simpler device, OMT, which separates an input linearly polarized wave into two components of the electric field parallel and perpendicular to the device symmetry axis at two output ports \\citep{wol02,men03}.  Traditionally, the septum polarizer has been known to perform well only within a relatively narrow frequency range, limited by the appearance of resonances.  On the other hand, the OMT has an advantage of being able to cover a wide bandwidth, and has been installed in modern telescopes, such as Atacama Large Millimeter/sub-millimeter Array.  A further comparison for the two kinds of devices reveals that the septum polarizer is good for measurements of Stokes $Q(\\equiv [\\langle E^{x}E^{x*}\\rangle-\\langle E^{y}E^{y*}\\rangle ]/2)$ and Stokes $U(\\equiv [\\langle E^{x}E^{y*} \\rangle+ \\langle E^{y}E^{x*}\\rangle ]/2)$, or the linear polarization, and the OMT good for those of Stokes $U$ and Stokes $V(\\equiv i[\\langle E^{x}E^{y*} \\rangle - \\langle E^{y}E^{x*}\\rangle ]/2)$, or the circular polarization, where $\\langle ...\\rangle$ are the time average. The reasons are as follows. Consider linearly polarized signals and denote the outputs of the septum polarizer to be the right-hand polarization electric field, $E^R\\equiv (E^x+iE^y)/\\sqrt{2}$ and the left-hand polarization electric field $E^L\\equiv(E^x-iE^y)/\\sqrt{2}$.  By cross-correlating $E^R$ and $E^L$, we obtain Stokes $Q$ as $\\langle E^R E^{L*}+E^RE^{L*}\\rangle$ and Stokes $U$ as $\\langle E^RE^{L*}-E^RE^{L*}\\rangle$.  In practice, one normally needs to amplify the weak incoming signals, with gains $G_R$ and $G_L$, immediately following the polarizer. The constructed Stokes parameters are in effect $\\langle G_R G_L\\rangle Q$ and $\\langle G_R G_L\\rangle U$, assuming the gains are real.  (See Section (6) for a discussion of complex gains.)  If one has an approximate knowledge of the average gains $G_R$ and $G_L$, the Stokes $Q$ and $U$ can be determined to an acceptable accuracy.  This is what can be achieved with a septum polarizer that measures $E^R$ and $E^L$ directly. On the other hand, if one adopts the OMT, the Stokes $U$ becomes $\\langle G_x G_y\\rangle[\\langle E^x E^{y*} \\rangle+\\langle E^y E^{x*} \\rangle]$ and the Stokes $Q$ becomes $\\langle G_x^2\\rangle \\langle|E^x|^2\\rangle-\\langle G_y^2\\rangle \\langle|E^y|^2\\rangle$.  While Stokes $U$ can be recovered in a similar manner as that with a septum polarizer, the recovery of Stokes $Q$ has a serious problem.  This is because not only the polarized signal is already mixed with the much stronger unpolarized sky, but the amplifiers also introduce substantial unpolarized noise to the signal.  Hence $|E^x|^2$ and $|E^y|^2$ contain almost the unpolarized radiation, and the recovery of weak polarized signals is reminiscent of the determination of a very small number by subtracting two big numbers from each other, for which any small error in the two big numbers will render a poorly determined small number.  The recovery of Stokes $Q$ is therefore only possible if the amplifier gains $\\langle G_x^2\\rangle$ and $\\langle G_y^2\\rangle$ can be calibrated to a high accuracy.  However, due to the presence of gain fluctuations in amplifiers, a telescope equipped with an OMT is often difficult to yield a well-determined Stokes $Q$.   On the other hand, when the polarized signal contains only Stokes $U$ and $V$, a similar argument applies, except for replacing $E^x$ by $E^R$, $E^y$ by $E^L$, and the OMT by the septum polarizer.  But, there have rarely been pure circular polarization signals in astronomical observations; hence the OMT is normally disfavored for astronomical polarization measurements and used mostly for the measurements of Stokes $I$. (Nevertheless, a sophisticated  solution for polarization measurements with the OMT has been proposed \\citep{men03}.) In spite of the great advantage of septum polarizers, they are not widely used in modern telescopes, simply because the polarizer has long been regarded as a narrow-band device. Therefore, It will be a great leap forward in the polarimeter instrumentation if this major weakness of the septum polarizer can be removed.  In this paper, we report our work specifically to address this issue.  This paper is organized as follows.  We introduce prior works on the septum polarizer in Sec.(2) and highlight the novel approach of the present work.  Sec. (3) outlines our design principles.  In Sec.(4), we report the measurement results of a polarizer fabricated for test. We convert the measurement results to the mutual leakage of Stokes parameters in Sec.(5). The leading-order calibration for reducing the Stokes $I$ leakage to other Stokes parameters is described in Sec.(6).  The conclusion is given in Sec.(7). ",
        "conclusions": "In this paper, we report a novel design of the septum polarizer that has a $42\\%$ well-performed bandwidth, from 77 GHz to 118 GHz, as opposed to the conventional notion of about $20\\%$ maximum bandwidth. The conventionally alleged maximum bandwidth was derived primarily from a consideration of a limit set by the cutoff frequency of the fundamental modes and the excitation frequency of high-order modes. Our polarizer, adopting a circular waveguide and housing a 5-step septum, is able to break the upper bandwidth limit and extends the usable frequency into the range where high-order $TM_{01}$ modes are excited.  This is made possible because we have uncovered an under-explored regime in which $TM_{01}$ excitations can be severely suppressed, and $TM_{01}$ resonances be entirely eliminated.   Particularly near the $TM_{01}$ excitation frequency, the septum can manage to avoid multiple reflections of the long wavelength modes inside the polarizer. In addition, the mutual leakage among all four Stokes parameters has been measured. It shows that this septum polarizer performs well, having $I$ to $Q$, $U$ leakage less than $2\\%$ and $Q-U$ mutual leakage less than $\\pm 1\\%$ in almost all frequencies.  A few dozen of polarizers of our design have been fabricated by conventional precision machining with $\\pm 5$ $\\mu m$ tolerance, and most of them have similar performances as the two modules reported here.  Due to the simplicity of the polarizer without any complicated component to assist widening the bandwidth, this polarizer sets a milestone for the instrumentation of polarization measurements.  This septum polarizers will be installed in an upgraded National Taiwan University (NTU)-Array prototype, where each receiver is equipped with 19 pixels of coherent detectors.  The $80-116$ GHz signals collected by the upgraded NTU-Array are to be processed by digital correlators that are designed to cover simultaneously the 36 GHz bandwidth for all pixels in all receivers with frequency resolution down to 100 kHz using software correlation.  In the context of this work, the backend digital processing capability of the telescope can help calibrate these polarizers with fine frequency resolution."
    },
    "1402/1402.4976_arXiv.txt": {
        "abstract": "In this letter we examine the evolution of the radial metallicity gradient induced by secular processes, in the disk of an $N$-body Milky Way-like galaxy. We assign a [Fe/H] value to each particle of the simulation according to an initial, cosmologically motivated, radial chemical distribution  and let the disk dynamically evolve for $\\sim 6$ Gyr. This direct approach allows us to take into account only the effects of dynamical evolution and to gauge how and to what extent they  affect the initial chemical conditions.  The initial [Fe/H] distribution increases with $R$ in the inner disk up to $R \\approx 10$ kpc and decreases for larger $R$. We find that the initial chemical profile does not undergo major transformations after $\\sim 6$ Gyr of dynamical evolution. The final radial chemical gradients predicted by the model in the solar neighborhood are positive and of the same order of those recently  observed  in the Milky Way thick disk.  We conclude that: \\item(1) the spatial chemical imprint at the time of disk formation is not washed out by secular dynamical processes, and \\item(2) the observed radial gradient may be the dynamical relic of a thick disk originated from a stellar population showing a positive chemical radial gradient in the inner regions. ",
        "introduction": "Observations suggest that the negative radial metallicity gradient, $d$[Fe/H]$/dR$, measured in the solar neighborhood close to the plane, flattens or even attains positive values if we select stars at higher $|z|$ from the galactic plane or with high [$\\alpha$/Fe] abundances (Allende Prieto et al. 2006; Boeche et al. 2013, Anders et al. 2013; Hayden et al. 2013). This effect can be a signature of a positive radial gradient in the thick disk population, as derived by Marsakov \\& Borkova (2005) and by Carrell et al. (2012), who analyzed kinematically selected thick disk samples. Alternatively, a similar feature can be simply produced by the variation of the fraction of thin disk and thick disk stars due to their different scale-height and scale-length (Anders et al. 2013). This scenario is indeed consistent with the short scale-length, $h_R\\sim 2$ kpc, recently estimated for $\\alpha$-enhanced stars by Bensby et al. (2011) and Cheng et al. (2012).  Furthermore, there is observational evidence indicating that the Milky Way (MW) is a barred galaxy. Non-axisymmetric structures in the disk (such as bars and spiral arms) drive galaxy evolution: they redistribute angular momentum among the different galactic components and can enhance radial migration inside the disk. It is still matter of debate the amount of impact that  these secular  processes can have  on the formation of the thick disk and on its chemical and kinematical properties \\citep{roskar2008, loebman2011, brunetti11, minchev12, Kubryk13}.  This work is aimed at analyzing whether dynamical effects can be responsible for the non-negative radial metallicity gradients observed today in the MW thick disk,  and suggests that such gradient can be regarded as  a cosmological feature.  We use an $N$-body disk that develops a bar and whose final population displays a radial metallicity gradient; this is obtained from an initial axisymmetric disk with a positive radial metallicity gradient for $R<10$ kpc, which has dynamically evolved for 6.1 Gyr. This secular evolution model was recently proven \\citep{Curir2012} to reproduce the observed velocity-metallicity correlation in the thick disk \\citep{spagna2010, lee2011}. In this letter, we show that the positive slope of the chemical radial gradient in the MW thick disk observed today in the solar neighborhood might be a signature remaining from the primordial population of disk stars. ",
        "conclusions": "The crucial role of a positive  slope $d{\\rm [Fe/H]}/dR$ in the inner disk for the primordial chemical distribution to produce a positive rotation-metallicity correlation in the thick disk \\citep{spagna2010, lee2011} was already shown in \\cite{Curir2012}. In this work we have analyzed how the same cosmologically motivated chemical distribution allows us to constrain the influence of the dynamical evolution on the metallicity gradients observed today in the MW thick disk.  We have drawn the radial and vertical chemical distributions from the same simulated MW galaxy used in \\cite{Curir2012}, which includes a {\\it positive} radial chemical gradient in the inner disk.  At 6.1 Gyr, within the solar neighborhood ($7.0$ kpc $ < R < 10.5$ kpc),  we still obtain a positive slope for the radial chemical distribution in the inner region. This slope  decreases  in going from $|z| < 0.5$ kpc to the region $2.5$ kpc $< |z| <3$ kpc ($d{\\rm [Fe/H]}/dR = 0.016\\pm0.002$ dex kpc$^{-1}$ and $d{\\rm [Fe/H]}/dR = 0.011\\pm0.007$ dex kpc$^{-1}$, respectively). The radial chemical gradients predicted by the model are at most within a factor of 2 of those observed by \\cite{Carrell12}, with the Photometric Distance Method, and  by \\cite{Hayden13}. Our  final positive radial gradients are  a direct consequence of our initial conditions, that include a positive radial gradient in the inner disk.  As shown in Figure \\ref{fig:grad_evol}, the radial gradients  in the solar neighborhood are of the same order from 1.8 Gyr to 6.1 Gyr, independently of the layer in $|z|$ (although some redistribution appears to take place).  Such a stability seen between $2$ and $6$ Gyr is consistent with the stability of the rotation-metallicity correlation pointed out by \\cite{Curir2012} in the same time range.  Furthermore, the vertical gradient in  the final configuration of our simulation is  very close to zero. This originates from the fact that we do not impose any dependence on $|z|$ in the initial chemical distribution. The lack of a $|z|$-dependence is also responsible from the lack of a trend for the radial gradients to increase with $|z|$.  We conclude that  the secular disk evolution does {\\it not} appear to be able to modify significantly the disk global chemical profiles, suggesting that the radial and vertical chemical distribution of the thick disk is likely to be a fossil signature of the original distribution.  Our results indicate that,   if the positive radial metallicity gradient in the solar neighborhood will be confirmed by the observations of  the  Milky Way  thick disk, this is  consistent with a thick disk population showing an  early positive [Fe/H] radial gradient in the inner disk ($R < 10$ kpc) and negative in the outer disk ($R > 10$ kpc). In the context of the inside-out formation scenarios of the galactic disk \\citep{spit01,Mott13}, such a gradient ``inversion'' derives from the strong infall of primordial gas that can occur at early times in the inner disk.    \\begin{figure*} \\includegraphics[angle=0,scale=0.8,trim=0 0 0 0]{fig3.ps} \\caption{Radial chemical distribution of the disk stars in a few kpcs of the MW solar circle, i.e. 7 kpc$<R< $ 10.5 kpc, after 6.1 Gyr of dynamical evolution. The different panels refers to different $|z|$ ranges. The black dashed lines are the best fit to the median values of [Fe/H] in each bin in R (of 0.05 kpc).} \\label{fig:metdist} \\end{figure*}  \\begin{figure*} \\includegraphics[angle=0,scale=0.8,trim=0 0 0 0]{fig4.ps} \\caption{Radial metallicity gradients of the disk stars with $7$ kpc $< R < 10.5$ kpc. The  squares are our estimates of the metallicity gradients at 6.1 Gyr from the dynamical $N$-body simulation. The abscissa of each point are the mean height  of the particles in the respective bin (broad $0.5$ kpc in $R$ or $|z|$). The error bars on the vertical axis are 3-$\\sigma$ width of the distribution, as 1-$\\sigma$ errors would be hardly appreciated in this figure. Circles and diamonds are the gradients  for $|z| > 1.5$ kpc measured by \\cite{Carrell12} through the Isochrone Distance and the Photometric Distance, respectively.} \\label{fig:gradients} \\end{figure*}   \\begin{figure*} \\includegraphics[angle=0,scale=0.8,trim=0 0 0 0]{fig5.ps} \\caption{Evolution of the radial gradient in the solar annulus in several layers in $|z|$, namely $0$ kpc $<|z|< 0.5$ kpc (empty circles), $0.5$ kpc $<|z|< 1.0$ kpc (filled circles), $1.0$ kpc $<|z|< 1.5$ kpc (diamonds), $1.5$ kpc $<|z|< 2.0$ kpc (triangles), $2.0$ kpc $<|z|< 2.5$ kpc (squares), $2.5$ kpc $<|z|< 3.0$ kpc (stars).} \\label{fig:grad_evol} \\end{figure*}"
    },
    "1402/1402.3551_arXiv.txt": {
        "abstract": "{} {We derive the stellar rotation curve of the Galaxy in the range of Galactocentric radii of $R=4-16$ kpc at different vertical heights from the Galactic plane of $z$ between -2 and +2 kpc. With this we reach high Galactocentric distances in which the kinematics is poorly known due mainly to uncertainties in the distances to the sources.} {We used the PPMXL survey, which contains the USNO-B1 proper motions catalog cross--correlated with the astrometry and near-infrared photometry of the 2MASS Point Source Catalog. To improve the accuracy of the proper motions, we calculated the average proper motions of quasars to know their systematic shift from zero in this PPMXL survey, and we applied the corresponding correction to the proper motions of the whole survey, which reduces the systematic error. We selected from the color-magnitude diagram $K$ vs. $(J-K)$ the standard candles corresponding to red clump giants and used the information of their proper motions to build a map of the rotation speed of our Galaxy.} {We obtain an almost flat rotation curve with a slight decrease for higher values of $R$ or $|z|$. The most puzzling result is obtained for the farthest removed and most off-plane regions, that is, at $R\\approx 16$ kpc and $|z|\\approx 2$ kpc, where a significant deviation from a null average proper motion ($\\sim 4$ mas/yr) in the Galactic longitude direction for the anticenter regions can be directly translated into a rotation speed much lower than at the solar Galactocentric radius. In particular, we obtain an average speed of $82\\pm 5$(stat.)$\\pm 58$(syst.) km/s (assuming a solar Galactocentric distance of 8 kpc, and a circular/azimuthal velocity of 250 km/s for the Sun and of 238 km/s for the Local Standard of Rest), where the high systematic error bar is due mainly to the highest possible contamination of non-red clump giants and the proper motion systematic uncertainty.} {A scenario with a rotation speed lower than 150 km/s in these farthest removed and most off-plane regions of our explored zone is intriguing, and invites one to reconsider different possibilities for the dark matter distribution. However, given the high systematic errors, we cannot conclude about this. Hence, more measurements of the proper motions at high $R$ and $|z|$ are necessary to validate the exotic scenario that would arise if this low speed were confirmed.} ",
        "introduction": "The rotation curve of the Milky Way disk was obtained using many different tracers and within different Galactocentric distance ranges. Dias \\& L\\'epine (2005) determined the rotation speed of the spiral pattern of the Galaxy by direct observation of the birthplaces of open clusters of stars in the Galactic disk as a function of their age, confirming that the spiral arms rotate like a rigid body, as predicted by the classical theory of spiral waves. Bobylev et al. (2008) used space velocities of young open star clusters and the radial velocities of HI clouds and star-forming regions to derive the Galactic rotation curve in the range of $3<R($kpc$)<12$; a more recent study by Reid et al. (2014) used proper motions of in-plane star-forming regions in radio and reached $R\\approx 16$ kpc. Bovy et al. (2012) measured a flat rotation curve in the range $4<R($kpc$)<14$ by fitting the radial velocities of some stars. Compilations of other rotation curve measurements are given by Sofue et al. (2009) and Bhattacharjee et al. (2013) and references therein.  In general, the outer rotation curve of the stellar population is more poorly determined mainly because of the inaccuracy in the distance determination to the stars or other sources (Sofue 2011), and we wish to improve this here. We use the red clump giant (RCG) population, which is a very appropriate standard candle (e.g., Castellani et al. 1992) with relatively low error bars on its distance estimate, and can be observed up to distances of $\\approx 8$ kpc, which allows one to reach $R\\lesssim 16$ kpc (L\\'opez-Corredoira et al. 2002). Williams et al. (2013) have used RCGs to derive this kinematical information, but only within the range  $6<R($kpc$)<10$, $|z|<2$ kpc, which does not reach the farthest regions of the stellar disk we aim to explore here. To derive the Milky Way rotation curve from proper motions of RCGs in the range $4<R($kpc$)<16$ for $|z|<2$ kpc.  The tool we used for our analysis of kinematics is a catalog of proper motions. Bond et al. (2010) have also made an extended analysis of the stellar kinematics, including the rotation speed, from proper motions with a revised version of the catalog SDSS-DR7. Aihara et al. (2011) reported later that there were significant problems with the measurement of the astrometry of SDSS-DR8 and previous versions, which in turn leads to problems with their proper motions as well, but Bond et al. (2010) may have taken them into account by analyzing the proper motion of quasars. Nonetheless, this analysis of the SDSS-DR7 survey does not cover the Galactic plane at $|b|<20^\\circ $, therefore it cannot be used to explore high values of $R$. Reid et al. (2014), as said, also analyze proper motions, but only for in-plane regions. Here we concentrate on the region $|b|<20^\\circ $.  The structure of this paper is as follows: in \\S \\ref{.data} the data source is described; the method to subtract the RCG stars is detailed in \\S \\ref{.redclump}; the rotation curve from proper motions is calculated in \\S \\ref{.proper}, with an extra section on how to correct for the systematic errors of proper motions in \\S \\ref{.syst}; this leads to the results in \\S \\ref{.rotation} and their dynamical consequences in \\S \\ref{.dyn}. A summary and conclusions are presented in the last section. Unless specified otherwise we assumed for the Sun a Galactocentric distance of 8 kpc and a circular/azimuthal velocity of 250 km/s, and a circular/azimuthal velocity of the Local Standard of Rest (LSR) of 238 km/s. ",
        "conclusions": "\\label{.concl}  Using only proper motions of sources identified as RCGs, we have derived the stellar rotation curve of the Galaxy in the range of Galactocentric radii of $R=4-16$ kpc for different vertical shifts from the Galactic plane of $z$ between -2 and +2 kpc. We obtained an almost flat rotation curve with a slight fall-off for higher values of $R$ or $|z|$. The most puzzling result is the observed low average rotation speed at the outest and most off-plane cases, that is, at $R\\approx 16$ kpc and $|z|\\approx 2$, with less than the half of the rotation speed with respect to its value in the solar neighborhood region: on average for $z=\\pm 2$ kpc, $V_c=82\\pm 5$(stat.)$\\pm 58$(syst.) km/s (assuming the systematic errors at $z=-2$ kpc and $z=+2$ kpc are independent, so they sum quadratically). If this speed is lower than 150 km/s (for $V_{g,\\odot }=250$ km/s), it would have important consequences for the analysis of the dark matter distribution in the outer part of our Galaxy. Nevertheless, at present we cannot exclude that the strong vertical gradient in the rotation speed is caused only by a combined systematic error of proper motions and contamination with non-RCGs. The assumptions made here, such as the approximation of circular orbits, might also be relaxed in the outest part of the disk. More measurements of the proper motions at high $R$ and $z$ are warranted to confirm or refute these results."
    },
    "1402/1402.6697_arXiv.txt": {
        "abstract": "Candidate galaxies at redshifts of $z \\sim 10$ are now being found in extremely deep surveys, probing very small areas. As a consequence, candidates are very faint, making spectroscopic confirmation practically impossible. In order to overcome such limitations, we have undertaken the CF-HiZELS survey, which is a large area, medium depth near infrared narrow-band survey targeted at $z=8.8$ Lyman-$\\alpha$ (Ly$\\alpha$) emitters (LAEs) and covering 10 deg$^2$ in part of the SSA22 field with the Canada-France-Hawaii Telescope. We surveyed a comoving volume of $4.7\\times 10^6$ Mpc$^3$ to a Ly$\\alpha$ luminosity limit of $6.3\\times10^{43}$ erg s$^{-1}$. We look for Ly$\\alpha$ candidates by applying the following criteria: i) clear emission line source, ii) no optical detections ($ugriz$ from CFHTLS), iii) no visible detection in the optical stack ($ugriz > 27$), iv) visually checked reliable NB$_J$ and $J$ detections and v) $J-K \\leq 0$. We compute photometric redshifts and remove a significant amount of dusty lower redshift line-emitters at $z \\sim 1.4 $ or $2.2$. A total of 13 Ly$\\alpha$ candidates were found, of which two are marked as strong candidates, but the majority have very weak constraints on their SEDs. Using follow-up observations with SINFONI/VLT we are able to exclude the most robust candidates as Ly$\\alpha$ emitters. We put a strong constraint on the Ly$\\alpha$ luminosity function at $z \\sim 9$ and make realistic predictions for ongoing and future surveys. Our results show that surveys for the highest redshift LAEs are susceptible of multiple contaminations and that spectroscopic follow-up is absolutely necessary. ",
        "introduction": "Finding the first stars and galaxies is one of the most important tasks to test our understanding of galaxy formation in the early Universe. The current theoretical models of when and how these first galaxies were formed can only be tested and improved by reliable detections of galaxies at the highest redshifts. The confirmation of galaxies at a redshift of $z\\sim 9-10$ would also allow the study of the epoch of reionization of the Universe. Measurements of the cosmic microwave background place this epoch at $z \\sim 10.6$ \\citep{Komatsu2011}, while \\cite{Fan2006} located the end of the reionization epoch at a redshift of at least $z\\sim 6$ by studying spectra of quasars at high redshift, where they found a lower limit to the neutral fraction of $\\sim 10^{-3} - 10^{-2}$.  A widely used technique to detect very distant galaxies is the Lyman break technique (LBG), pioneered by \\cite{Steidel1996} (see also \\citealt{Guhathakurta1990}), which looks at a distinctive break in the UV spectrum of star-forming galaxies. More generally, one can use deep data in several broadbands to derive a redshift-probability distribution by fitting spectral energy distributions (SED) based on galaxy templates \\citep[e.g.][]{McLure2011}.  Using the Lyman Break method, candidate galaxies have been found at very high redshifts \\citep[$z \\sim 7$, e.g.][]{Bouwens2011z7,Finkelstein2012,Oesch2012z8,McLure2012} and even $z\\sim10$ \\citep[]{Ellis2013,Oesch2013,Bouwens2013}, but the great majority of these are too faint to confirm spectroscopically. \\cite{Lehnert2010} claimed the spectroscopic detection of a $z=8.6$ Lyman-$\\alpha$ (Ly$\\alpha$) line of a LBG in the Hubble Ultra Deep Field. However, \\cite{Bunker2013} were unable to reproduce the detection with two independent sets of observations, leading to the suggestion that it could be an artefact. \\cite{Brammer2013} found a tentative emission line that could be Ly$\\alpha$ at $z = 12.12$ using the HST WFC3 grism, but this is only a $<3 \\sigma$ detection and could be a lower redshift interloper. Recently, \\cite{Finkelstein2013} report the detection of a Ly$\\alpha$ emission line in a $z=7.51$ LBG, although the line is very close to a sky-line making identification significantly more difficult. Other attempts have been made, but so far no $z>7.5$ galaxy has been spectroscopically confirmed. There is a spectroscopic redshift determination of a $z=8.2$ Gamma Ray Burst \\citep{Tanvir2009}, but not for its host.  Another successful technique to detect very high redshift ($z \\sim 4- 7$) galaxies is the narrow-band technique, which targets Lyman-$\\alpha$ emitters (LAEs; e.g. \\citealt{Pritchet1994,Thompson1994,Thompson1995,Hu1996,Cowie1998,Hu1998,Thommes1998,Rhoads2000,Rhoads2003,Rhoads2004,Fynbo2001,Hu2002,MalhotraRhoads2002,MalhotraRhoads2004,Fynbo2003,Ouchi2003,Hu2004,Taniguchi2005,Iye2006,Kashikawa2006,Shimasaku2006,Ouchi2008,Finkelstein2009,Ota2010,Hibon2011}).  Using the narrow-band technique one can search for sources with emission lines at specific redshifts, by looking at the excess the narrow-band has over the broadband. This way sources for which the continuum is too faint to be detected, can still be identified due to the bright emission lines. However, most emission line galaxies detected in narrow-band surveys are lower redshift interlopers such as H$\\alpha$ and [O{\\sc ii}] \\cite[e.g.][]{Sobral2012}, which have to be identified using multiwavelength observations. Because the narrow-band is only sensitive to sources emitting in a small range of wavelengths, they can be used to look at a slice of redshifts and therefore a well-known comoving volume. Moreover, spectroscopic follow-up of high-redshift candidates is a priori easier for candidates detected by the narrow-band technique, as these candidates will have strong emission lines. Currently, the most distant spectroscopically confirmed NB-selected LAE is at a redshift of 6.96  \\citep{Iye2006}, which is detected with narrow-band imaging from the Subaru telescope.  \\begin{figure} \\centering \\includegraphics[width=8cm]{./REPORT_FIGURES/transmission.png} \\caption{\\small{The atmospheric transmission in the near-infrared $J$ band, normalised to the maximum transmission for each curve. Atmospheric data is from Mauna Kea, Gemini Observatory \\citep{Lord1992}, the airmass is 1.0 and water vapour column is 1.0 mm. The background (grey) atmospheric (OH) emission lines are also shown. The NB$_J$ filter used in this paper is transparent at wavelengths where there are no strong OH lines and also at wavelengths where the atmosphere is at its maximum transparency, thus allowing us to obtain deep observations in relatively little time.}} \\label{fig:filters} \\end{figure}  To observe even higher redshift galaxies, observations in the near-infrared are required. Unfortunately, at these wavelengths there is significant foreground emission due to OH molecules in the Earth's atmosphere. Some OH windows exist at wavelengths where the atmosphere is transparant to radiation. It is possible to observe using narrow-band filters in these windows very effectively and several filters have been developed for this purpose (see Figure $\\ref{fig:filters}$). Recent studies led to the identification of candidate Ly$\\alpha$ emitters at a redshift of $z = 7.7$, but none of these has been spectroscopically confirmed yet \\citep[]{Hibon2010,Tilvi2010,Clement2012,Krug2012,Jiang2013b}.  Some attempts at somewhat higher redshifts ($z\\sim9$) were made to detect Ly$\\alpha$ \\citep[]{WillisCourbin2005,Cuby2007,Willis2008,Sobral2009b}. The properties of such galaxies would provide strong tests of current models of galaxy formation and evolution and even the confirmation of just one luminous Ly$\\alpha$ emitter at this redshift will be suitable for the study of these sources way before the next generation of telescopes, such as \\emph{JWST} or the E-ELT.  Ly$\\alpha$ radiation is much more attenuated by a neutral intergalactic medium (IGM) than an ionized IGM, so large samples of Ly$\\alpha$ emitters at these redshifts could be used to derive properties of the IGM at these early times.  Current simulations \\citep[e.g][]{Iliev2008} suggest that reionization started at the most overdense regions in the Universe, where ionizing sources nurtured expanding shells of ionized gas in the IGM. As Ly$\\alpha$ radiation is easily absorbed by a neutral medium \\citep{MalhotraRhoads2004}, Ly$\\alpha$ emitters can only be observed once the ionized zone around them is large enough for the Ly$\\alpha$ radiation to escape. This is expected to lead to a negative evolution in the Ly$\\alpha$ luminosity function and dropping escape fraction of Ly$\\alpha$ radiation at higher redshifts. Considerable effort has been put in spectroscopically studying the evolution of the Ly$\\alpha$ line in Lyman break galaxies (LBGs) at high redshifts \\citep[e.g.][]{Fontana2010,Pentericci2011,Vanzella2011,Ono2012,Schenker2012,Caruana2013,Finkelstein2013}. Recent non-confirmations and low success-rates at $z>7$ for their spectroscopic confirmation are interpreted as a signature that reionization is not yet completed at these redshifts. \\cite{Treu2013}, for example, find that at $z\\sim8$ Ly$\\alpha$ emission of LBGs is suppressed by at least a factor of three.  For LAEs, it is found that up to at least a redshift of $z\\sim6$ the Ly$\\alpha$ luminosity function is remarkably constant \\citep[e.g.][]{Shimasaku2006,Hu2004,Ouchi2008}. This indicates that LAEs are relatively more common and more luminous at earlier epochs, compared to LBGs (as the UV LF drops quickly in this redshift range; \\citealt{Bouwens2007}). At $z\\sim6-8$ there is evidence for evolution of the characteristic luminosity, but these samples, including failed attempts at $z=7.7$, can be significantly affected by cosmic variance, probing $\\lsim 1$ deg$^2$ \\citep[e.g.][]{Ouchi2010,Clement2012}.  At the bright end, however, the evolution could plausibly be very different. Luminous sources can ionize their own surroundings to allow Ly$\\alpha$ photons to escape, as they redshift out of restframe-resonance wavelength in about 1 Mpc \\citep[e.g.][]{Barton2004,CenHaiman2000,Curtis-Lake2012}. Furthermore, the observed clustering of Ly$\\alpha$ emitters is expected to increase at higher redshift, as neighbouring sources will have larger overlapping ionized spheres and therefore a higher fraction of escaped Ly$\\alpha$ photons \\citep[e.g.][]{Ouchi2010}.  In order to find the most luminous Ly$\\alpha$ emitters in the epoch of reionization which would be suitable for spectroscopic follow-up, we have undertaken the widest area search with a near infrared narrow-band filter to date. This paper is organised in the following way. \\S2 presents the details of the observations, and describes the data reduction, calibrations and source extraction. \\S3 presents the criteria for sources being selected as Ly$\\alpha$ candidates and the results from the narrow-band search. \\S4 presents the spectroscopic follow-up observations and results. \\S5 discusses the results such as constraints on the Ly$\\alpha$ $z=8.8$ luminosity function, and our survey is compared to past and future surveys. Finally, \\S6 outlines the conclusions. A H$_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$ cosmology is used  and all magnitudes are in the AB system, except if noted otherwise. ",
        "conclusions": "We have conducted a very wide narrow-band survey over 10 deg$^2$ in the near infrared and identified 6315 line-emitters using a $1.19 \\mu$m narrow-band filter. In this work we identified possible $z = 8.8$ Ly$\\alpha$ candidates in the sample of line-emitters and followed-up the strongest ones spectroscopically. The main conclusions are:  \\begin{itemize} \\item A significant fraction ($\\sim 300$) of the line-emitters are consistent with being at high redshift ($z > 3$), of which some might be Ly$\\alpha$ at $z = 8.76$. This narrow-band survey increased the probed volume by half an order of magnitude compared to previous surveys and is thus sensitive to the rarest and most luminous sources. \\item By doing careful visual checks of the robustness of the detections and by excluding line-emitters which are detected in any of the optical bands and which show a red $J-K$ colour, we find 13 possible Ly$\\alpha$ candidates. We order them in different groups based on their broadband photometric constraints. The two most robust candidates have reliable detections in narrow-band, strongest constraints from photometric redshifts, $iz-J$ break and robust $J$ detections. \\item 90\\% of the high redshift candidate line-emitters, selected on having no/very faint flux in the optical, are lower redshift interlopers. By including the $K$ band and computing photometric redshifts we find that approximately 40\\% are H$\\beta$/[O{\\sc iii}] at $z=1.4$, 25\\% [O{\\sc ii}] at $z=2.2$ and 15\\% are candidate AGN emission-lines (e.g. carbon or magnesium) at $z\\sim3-6$, while the remaining are likely very faint lower redshift sources like Pa$\\gamma$ at $z = 0.09$ or He{\\sc ii} at $z = 0.44$. \\item Spectroscopic follow-up of the two most robust Ly$\\alpha$ at $z=8.8$ candidates, two sources with the largest EWs and another with brightest J failed to confirm these sources as line-emitters. This is probably caused by a combination of spurious sources, variability and (although unlikely) solar system objects. This result has very strong implications to current and future candidates for LAEs at $z=7.7$ and $z>8$. \\item After the follow-up, we put the strongest constraints on the bright end of the luminosity function with half an order of magnitude improvement in the probed volume and it could still mean little to no evolution in the luminous end. \\item Using an optimistic upper-limit to the LF and a lower redshift extrapolation, we estimate the number of LAEs that will be detected by the completed VISTA NB118 GTO survey to be between $0.001$ and $1.19$ with the current depth and to be between $0.19$ and $23.35$ for the estimated final depth of the ongoing UltraVista NB118 survey. \\item Because of the lack of a comparably wide surveys, it is difficult to study the evolution of the bright end of the LF and extrapolations from other considerably smaller surveys at lower redshifts are unusable. Although the number density of Ly$\\alpha$ emitters is expected to decline at higher redshifts, this isn't necessary the case for the bright end of the luminosity function, because of the topology of reionization. It is therefore of utmost importance to study the bright end of the Ly$\\alpha$ luminosity function at lower redshifts in order to understand the evolution in the LF completely. \\item As our strongest candidates looked realistic in the images and had realistic physical properties based on the photometry, but still are not confirmed, we highlight the necessity for all other surveys to do this spectroscopic follow-up, especially when candidates are based on just a single-band detection. This has significant consequences for any similar and for deeper surveys, clearly pointing out that despite sources passing all tests, only spectroscopic observations can confirm them. \\end{itemize}"
    },
    "1402/1402.1762_arXiv.txt": {
        "abstract": "The interest on the solar diameter variations has been primary since the scientific revolution for different reasons: first the elliptical orbits found by Kepler in 1609 was confirmed in the case of the Earth, and after the intrinsic solar variability was inspected to explain the climate changes. The CCD Solar Astrolabe of the Observat\\'orio Nacional in Rio de Janeiro made daily measurements of the solar semi-diameter from 1998 to 2009, covering most of the cycle 23, and they are here presented with the aim to evidence the observed variations. Some instrumental effects parametrizations have been used to eliminate the biases appeared in morning/afternoon data reduction. The coherence of the measurements and the influence of atmospheric effects are presented, to discuss the reliability of the observed variations of the solar diameter. Their amplitude is compatible with other ground-based and satellite data recently published. ",
        "introduction": "The interest on the solar diameter variations has been prioritary during the last four centuries of astronomy with and without the telescope. During 1600 the verification of the Keplerian hypothesis of elliptical Earth's orbit was possible thanks to measurements of the seasonal variations of the angular diameter of the Sun made by G. D. Cassini and G. B. Riccioli from 1655 to 1666 at the {\\sl Heliometer} of S. Petronio in Bologna. This was a pinhole solar telescope\\footnote{The use of the word {\\sl telescope} for an optics-less instrument like a giant pinhole camera realized into a church provided by a meridian line accurately draft on the ground, may appear awkward. But the Greek word's meaning is related to the capability of seeing at far objects, and the implicit possibility to have magnified images, and this is possible with pinhole instruments.} of the Basilica of St. Petronius in Bologna \\cite{Manfredi}. Later in 1700 the interest of the astronomers was more focused on the Earth's axis orientation, with the variation of the obliquity $\\epsilon$. To measure this parameter it was necessary to know better the position of the solar center, than the diameter. The giant pinhole telescopes \\cite{Heilbron} could operate on a complete yearly range of zenithal distances of $\\sim 47^{o}=2\\times\\epsilon$ without optical distortions, excepted the atmospheric refraction, and for this reason these instruments continued successfully their operations until the end of the century  even if their imaging capability was comparatively lower than usual  telescopes. The heliometer as we know nowadays was invented in 1743 by S. Savery \\cite{Short1753} and built by J. Dollond ten years later, but the solar diameter has been measured with timing meridian transits at the Greenwich Observatory, starting with N. Maskelyne from 1765 to 1810, and the influence of a variable {\\sl personal equation} with the age, either in reaction time either in contrast sensitivity, has been recognized as the most reliable cause of the recorded variations.  The heliometric method, which does not deal with reaction times, but only with the personal contrast sensitivity, was considered more accurate than \\cite{Auwers1890}. The nineteenth century was the golden age of the solar astrometry, and the heliometers built by J. Fraunhofer were so accurate that were used by F. W. Bessel for the measurement of the first stellar parallax \\cite{Bessel1838}. The measurements of the solar diameter variations became a routine with specific instrument designed for, and the observatories involved in the meridian transit measurements were also Neuchatel, Oxford, Washington \\cite{Auwers1890}, and Rome-Campidoglio \\cite{Gething1955}. The current standard value of the solar semi-diameter has been fixed to 959.63 arcsec\\footnote{This angular dimension corresponds to 696000 Km at 1 AU.} by A. Auwers in 1891 \\cite{Auwers1891}. The words of A. Secchi in the book \"Le Soleil\" expressed clearly the status of art at the fall of the century \\cite{Secchi1875}: {\\sl What strikes even more is to see that, despite the variety of methods and instrumental perfection, the measurement of the solar diameter made very little progress.}\\footnote{Ce qui frappe encore davantage, c'est de voir que, malgr\\'e la vari\\'et\\'e des m\\'ethodes et la perfection des instruments, la mesure du diam\\`etre solaire a fait bien peu de progr\\`es (Secchi, 1875 p. 210).} ",
        "conclusions": "The aforementioned words of A. Secchi in 1875 could be applied also nowadays, even if meanwhile the errorbars of these measurements have been reduced of more than a factor of 10-100. The influence of atmospheric turbulence and the environmental parameters can affect a single measurements with perturbations bigger than the real variations of the solar diameter. The era of satellites opened the way to the milliarcsecond solar astrometry, but for the knowledge of the solar diameter fluctuations on timescales longer than the solar cycle, their lifetime is very short and their operational cost bigger than astrolabes and heliometers.  The Solar Heliospheric Observatory SOHO has the longer operating time until now, but its instruments were not designed for astrometry. {\\it Nonetheless its data have been exploited to show that the solar diameter has not changed within a few milliarcsecond per year...} (Delmas et al. (2006) p. 1567 p. 4.)  Because of the nature of planetary transit measurements of solar diameter, an excellent angular resolution is achievable by a good timing resolution. Hence the Mercury transits of 2003 and 2006 observed by SOHO provided the occasion to measure the solar diameter based only on timing accuracy  independently on angular references \\cite{Emilio12}. Therefore these measurements made with SOHO in 2003 and 2006, 960.08 arcsec, can be fairly compared with the analogous measurements made by PICARD satellite during the transit of Venus of 2012 in the same waveband, 959.86 arcsec, concluding that the variations of the diameter from 2003 to 2012 has been $0.22\\pm0.10$ arcsec. This is in perfect agreement with the variations measured by the Rio Astrolabe during the cycle 23, and this may explain directly the scatter between the yearly averages found in the aforementioned series of Greenwhich, Campidoglio, Rio de Janeiro and Calern.  The astrolabe of Rio de Janeiro allowed to monitor the solar during more than a decade without the costs of a space mission and its limited operating lifetime. Future results from Hinode, SDO and PICARD satellite can better elighten this field of research, and instruments like PICARD-Sol and the Reflecting Heliometer of Rio de Janeiro will help to maintain the 0.01 arcsec of accuracy in solar astrometry typical of the satellites to ground-based observations, to prolong the series on the solar diameter for the years to come.  The measurements of the solar diameter made with the solar Astrolabes have been debated since many decades. In this paper all original data of the observations made in Rio de Janeiro over the whole cycle 23 have been here presented as appendix from \\cite{Boscardin2011} and discussed. The scope was also to share with the international heliophysics community these data with the problems of their interpretations. Since this kind of problems are common with all ground-based observations \\cite{Gething1955} their final solution will take full advantage also of these data, to verify, for example, the hypotheses on the possibility that the solar activity could influence indirectly the measurements as some scholars claimed \\cite{Rozelot2003}.   \\begin{acks} C.S. thanks J.P. Rozelot and W. Dziembowski for fruitful discussions and suggestions. \\end{acks}"
    },
    "1402/1402.2282_arXiv.txt": {
        "abstract": "{ Comprehensive studies of Wolf-Rayet stars were performed in the past for the Galactic and the LMC population. The results revealed significant differences, but also unexpected similarities between the WR populations of these different galaxies. Analyzing the WR stars in M31 will extend our understanding of these objects in different galactic environments. }{ The present study aims at the late-type WN stars in M31. The stellar and wind parameters will tell about the formation of WR stars in other galaxies with different metallicity and star formation histories. The obtained parameters will provide constraints to the evolution of massive stars in the environment of M31. }{ We used the latest version of the Potsdam Wolf-Rayet model atmosphere code to analyze the stars via fitting optical spectra and photometric data. To account for the relatively low temperatures of the late WN10 and WN11 subtypes, our WN models have been extended into this temperature regime. }{ Stellar and atmospheric parameters are derived for all known late-type WN stars in M31 with available spectra. All of these stars still have hydrogen in their outer envelopes, some of them up to 50\\% by mass. The stars are located on the cool side of the zero age main sequence in the Hertzsprung-Russell diagram, while their luminosities range from $10^{5}$ to $10^{6}\\,L_{\\odot}$. It is remarkable that no star exceeds $10^{6}\\,L_{\\odot}$. }{ If formed via single-star evolution, the late-type WN stars in M31 stem from an initial mass range between $20$ and $60\\,M_{\\odot}$. From the very late-type WN9-11 stars, only one star is located in the S Doradus instability strip. We do not find any late-type WN stars with the high luminosities known in the Milky Way. } ",
        "introduction": "\\label{sec:intro} M31 or the Andromeda galaxy is the largest member of the Local Group, which contains only two other spiral-type galaxies, M33 and the Milky Way. Because of the very low foreground extinction towards M31 and the known distance, this galaxy is ideal for studying bright, resolvable stellar objects, such as Wolf-Rayet (WR) stars. Analyzing the WR stars in M31 will extend our knowledge about the formation and evolution of massive stars in other galaxies and allow for a comparative analysis of their WR populations.  Previous studies chiefly focused on the WR population in the Milky Way and the Magellanic Clouds. One of the reasons for this limitation was the paucity of available spectra of WR stars in other galaxies. The advent of multi-object spectroscopy has greatly facilitated such observations. Good quality spectra of extragalactic WR stars are now becoming available, allowing for qualitative analyses. A large set of WR spectra for M31 have been published by \\citet{NMG2012} and provided the basis for this work. Photometric data of extragalactic WR stars are contained in the Local Group Galaxy Survey (LGGS) by \\citet{Massey+2006}.  Spectroscopically, the WR stars are divided into the WN, WC, and WO subclass. The WN stars show prominent nitrogen emission lines in their spectra while the WC and WO stars have prominent carbon and oxygen emission lines. In contrast to the WC stars, the oxygen emission lines are significantly stronger in the WO star spectra, including a prominent \\ion{O}{vi} emission line at $3811$-$34\\,$\\AA. All subclasses are further split into a sequence of subtypes, defined by the equivalent width or peak ratio of certain emission lines \\citep{CHS1995, vdH2001}. For the WN stars, these are the subtypes WN2 to WN11. The WN2 to WN6 subtypes are also called ``early'' types, while the WN8 to WN11 are referred to as ``late'' subtypes. The WN7 subtype is something in between, although it is often included under the late subtypes. The WN9 to WN11 subtypes were introduced by \\citet{CHS1995} as a finer replacement of the former Ofpe/WN9 classification and are sometimes referred to as ``very late'' WN types (WNVL: very late WN).  In contrast to the early subtypes, late-type WN stars typically have significant amounts of hydrogen \\citep{HGL2006}. In the past years, hydrogen-rich WN stars, sometimes classified as WNh, have become an interesting topic of research. Studies such as \\citet{Graefener+2011} have pointed out that at least a significant fraction of these stars are not ``classical'' WR stars in the sense that they are stripped cores of evolved massive stars. Instead, WN stars of this kind are most likely core-hydrogen burning, and they form the most luminous group within the WR population. The most luminous WR stars in the Large Magellanic Cloud (LMC) and the Milky Way are in fact such hydrogen-rich WN stars \\citep{BOH2008,Crowther+2010}.  In this paper we analyze the late-type WN stars with subtypes ranging from WN7 to WN11. Apart from two stars with insufficient spectra, we cover the whole known sample of late-type WN stars in M31 with available optical spectra.  This present paper builds on our recent work that analyzes the WR star populations in the Milky Way \\citep{HGL2006,SHT2012} and the LMC \\citep{Hainich+2014}. Extending the analysis to galaxies beyond the Milky Way may in principle allow study of WR stars at higher metallicities. According to \\citet{ZKH1994}, the metallicity in M31 has a value of  $\\log(\\text{O/H}) + 12 = 8.93$, which is a bit higher than the solar value of $8.7$ \\citep{EP1995}. However, as in our Milky Way, the metallicity is not uniform across the whole galaxy but instead increases towards the central bulge. Figure\\,\\ref{fig:wnpos} shows the positions of the analyzed stars in M31. It is significant that none of the known late-type WN stars in M31 are located in the inner part of this galaxy. Instead the stars mostly reside in a zone between 9\\,kpc and 15\\,kpc from the center of M31. It is exactly this region that \\citet{vdB1964} found to be the most active in forming stars. Because of their non-central location it might be that the metallicity of our sample is not significantly higher than the one used in our Galactic WN models \\citep[cf.][]{HG2004}.  The distance to M31 is well known \\citep[$d\\,=\\,0.77\\,$Mpc,][]{CNG}. This offers a big advantage for analyzing of stellar populations. Among the Galactic stars, there is only a limited subsample for which distances can be inferred from cluster or association membership. Analyzing the M31 WN stars will therefore allow us to crosscheck our findings from the Galactic WR analyses and compare them with the LMC results where we have a different type of galaxy and lower metallicity.  In the next section, we briefly characterize the stellar wind models used in this work. In Sect.\\,\\ref{sec:sample} we report the results and compile the obtained stellar parameters. In Sect.\\,\\ref{sec:evol} we discuss the results and compare them with theoretical stellar evolution tracks. The conclusions are drawn in the Sect.\\,\\ref{sec:conclusions}. In the appendix (Sect.\\,\\ref{appsec:specfits}) we show all spectral fits obtained in this work.  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{m31wnpos}} \\caption{The positions of the analyzed WN stars in M31. Numbers correspond to the M31WR numbers as in Table\\,\\ref{tab:fitresults}. The background image is taken from Palomar STScI DSS. The size of the image is $\\approx 1.9^\\circ \\times 2^\\circ$. The gray dashed ellipse roughly corresponds to a radius of $9\\,$kpc around the center of M31, corresponding to the inner boundary of the region with the most star formation \\citep{vdB1964}.} \\label{fig:wnpos} \\end{figure} ",
        "conclusions": "  \\begin{itemize} \\item All stars in our sample have luminosities $L$ between $10^5$ and $10^6\\,L_{\\odot}$. Notably, we do not find any star in our sample that exceeds $10^6\\,L_{\\odot}$. \\item The absolute visual magnitude $M_{\\text{V}}$ shows a significant scatter, even within the same WN subtype. This sets limits to the applicability of a subtype magnitude calibration. \\item Since the current catalog of WN stars in M31 does not include the central region of this galaxy, we cannot rule out that this special environment hosts WN stars with luminosities $L > 10^6\\,L_{\\odot}$, especially since we know that such stars exist in the LMC and the Galactic center region. \\item The late-type WN stars in M31 are in good agreement with the latest set of the Geneva tracks with $Z = 0.014$. Only the Geneva tracks with rotation can reproduce the lowest luminosities in our sample. Otherwise, the different sets of tracks \\citep{Georgy+2012,EIT2008} can all explain the observed WNh stars because these stars are not yet in the final WR stages for which the tracks differ more. \\item If formed via single-star evolution, the analyzed late-type WN stars in M31 stem from an initial mass range between $20$ and $60\\,M_{\\odot}$. The lower limit might in fact be higher if the few stars with the lowest luminosities were formed via binary evolution. \\item Our spectral fits reveal that the analyzed late-type WN stars have similar chemical compositions to their Galactic counterparts. This is in line with the fact that the stars are not located in the inner part of M31, where we would expect higher metallicities. The obtained HRD positions of the analyzed stars do not provide additional constraints since they can be reproduced by evolutionary models with both Galactic and higher metallicity. \\item The number of very late WN stars (WN9-11) is low, which could partly be a selection effect of the detection method \\citep{NMG2012}. While a lot of additional WN9 stars have been discovered in the Milky Way in obscured regions with infrared observations, the number of currently known WN10 and WN11 stars in M31 is about the same as in the Milky Way and the LMC. \\item Only one object, the WN10 star \\#148, is located in the S\\,Dor instability strip in the HR-diagram. However, the HRD positions do not rule out that some of the stars of very late-type WN subtypes might be LBVs in the quiescent stage. From the luminosity range, it is hard to tell whether the WN stars have lost parts of their hydrogen via an LBV outbursts or in the red supergiant (RSG) phase. \\end{itemize}  In a forthcoming paper we will analyze the early WN subtypes to get a more complete picture of the WN population in M31."
    },
    "1402/1402.5694_arXiv.txt": {
        "abstract": "We analyzed 3 years of data from the \\kepler\\ space mission to derive rotation periods of main-sequence stars below 6500\\,K. Our automated autocorrelation-based method detected rotation periods between 0.2 and 70 days for 34,030 (25.6\\%) of the 133,030 main-sequence \\kepler\\ targets (excluding known eclipsing binaries and Kepler Objects of Interest), making this the largest sample of stellar rotation periods to date. In this paper we consider the detailed features of the now well-populated period-temperature distribution and demonstrate that the period bimodality, first seen by McQuillan, Aigrain \\& Mazeh (2013) in the M-dwarf sample, persists to higher masses, becoming less visible above 0.6\\,$M_\\odot$. We show that these results are globally consistent with the existing ground-based rotation-period data and find that the upper envelope of the period distribution is broadly consistent with a gyrochronological age of 4.5\\,Gyrs, based on the isochrones of \\citet{bar07}, \\cite{mam+08} and \\cite{mei+09}. We also performed a detailed comparison of our results to those of  \\cite{rei+13} and \\cite{nie+13}, who have measured rotation periods of field stars observed by \\kepler. We examined the amplitude of periodic variability for the stars with detected rotation periods, and found a typical range between $\\sim 950$\\,ppm ($5^{\\rm th}$ percentile) and $\\sim 22,700$\\,ppm ($95^{\\rm th}$ percentile), with a median of $\\sim 5,600$\\,ppm. We found typically higher amplitudes for shorter periods and lower effective temperatures, with an excess of low-amplitude stars above $\\sim 5400$\\,K. ",
        "introduction": "Throughout their main-sequence lifetime, stars lose angular momentum and spin down. For intermediate- and low-mass stars, this angular momentum loss is thought to occur via a magnetized wind that is linked to the stellar outer convection zone \\citep[e.g.,][]{kaw88, bou+97}. Therefore, the present-day rotation period reflects the integrated angular momentum loss history of the star.  It has been claimed that for main-sequence stars, the present day rotation period does not depend on the initial conditions, and therefore a tight relationship exists between the period, age and mass, which is known as gyrochronology, for which a range of empirical relations have been derived \\citep[e.g.,][]{kaw89, bar03, bar07, bar10, mam+08}. In recent years there has also been considerable work towards understanding and modeling the underlying physics behind gyrochronology \\cite[][and references therein]{rei+12, mat+12, gal+13}. Therefore, deriving the rotation period for a large number of main-sequence stars has been a long-standing goal in stellar astrophysics, with the potential to shed light on stellar, and even planetary system angular-momentum evolution on one hand, and provide a new method to probe Galactic star formation history and structure on the other hand.  In recent years, the study of stellar rotation has seen significant advances thanks to new observational programs and techniques. Previously, measurements of stellar rotation were typically performed using spectroscopy, via the rotational broadening of absorption lines \\cite[e.g.,][]{kal89}. However, this technique yielded only model-dependent constraints on the rotation rate due to the dependence on stellar radius and inclination, and were limited to relatively fast rotators.  The era of exoplanet transit surveys provided a new approach to stellar rotation studies, making it possible to measure the rotation period directly for a large number of stars simultaneously. This is done by detecting quasi-periodic brightness variations in the time series photometry, caused by magnetically active regions repeatedly crossing the visible hemisphere as the star rotates \\citep[e.g.,][]{irw+09a}.  Open cluster surveys over the past decade have provided thousands of rotation periods for low-mass stars with ages up to $\\sim 650$\\,Myr, an overview of which can be found in \\cite{irw+09} and \\cite{bou+13}. More recent additions to the literature on rotation of young low-mass stars include studies of M35 (310 periods) \\citep{mei+09},   Coma Berenices \\citep{col+09}, M34 (83 periods) \\citep{mei+11}, and h Persei (586) \\cite{mor+13}. Together, these data have provided a relatively complete, but complex picture of pre-main-sequence rotational evolution.  Until recently, rotation periods for older stars remained scarce, since their slow rotation and low-amplitude modulation make them difficult to detect with either spectroscopy or ground-based photometry. Notable exceptions include $112$ main sequence F, G and K stars observed as part of the Mount Wilson H-K project \\citep[][and references therein]{bal+96}, $1727$ mid-F to mid-K stars observed by the {\\it CoRoT}\\ satellite \\citep{aff+12}, 41 low-mass ($0.1$--$0.3\\,M_\\odot$) stars from the \\mearth\\ survey \\citep{irw+11}, and $\\sim 2000$ field K and M stars observed by the HATNet survey \\citep{har+11}. The periods of \\cite{aff+12} are based on up to 5 months of CoRoT data and yet include a significant fraction of periods greater than 60 days. Since many of these periods are more than half the length of the dataset, they should be treated with caution and we opted to omit these results from the comparison section of this paper.  Asteroseismology studies performed on a small subset of \\kepler\\ and CoRoT targets also provide an insight into stellar rotation, since the oscillation spectrum of a star can reveal information about surface and interior rotation rates \\citep[e.g.,][]{giz02, giz+04}. For a recent review on observational studies of stellar rotation, see \\cite{bou+13}.  Data from the Kepler space mission \\citep{bor+10, koc+10} are revolutionizing the study of stellar rotation, providing 4\\,yrs of almost uninterrupted photometry with an unprecedented level of precision and time sampling, for a very large sample of stars. This allows rotation periods to be detected for slowly rotating stars with low-amplitude modulation, for a well-defined large sample.  Rotation studies on subsets of the \\kepler\\ data include the measurement of 265 rotation periods for stars with $T_{\\rm eff} \\le 5200\\,$\\,K and $\\log g \\ge  4.0$\\,dex observed by \\kepler\\ for 1--2 quarters through the Cycle 1 Guest Observer program \\citep{har+12}, 1570 M-dwarf rotation periods using quarters 1--4 (Q1--4) \\citep{mcq+13a}, 737 of the Kepler Objects of Interest (KOIs) using Q3--14 data \\citep{mcq+13b}, $\\sim 950$ of the KOIs using Q9 \\citep{wal+13}.  Two previous studies focussing on the broader \\kepler\\ sample are those of \\cite{rei+13}, with an emphasis on differential rotation, who derive $\\sim 24,000$ periods using Q3, and \\cite{nie+13}, who measured $\\sim 12,000$ periods from Q2--Q9, and compare to previous spectroscopic studies. These studies, discussed in Section~\\ref{sec:disc}, use Fourier-based detection methods and conservative automatic selection criteria.  In this paper we describe the design and application of AutoACF, an automated version of the autocorrelation function (ACF) method introduced by \\cite{mcq+13a} for stellar rotation period measurement. Using AutoACF we derived 34,030 rotation periods from 133,030 main-sequence targets observed by  \\kepler. We outline the sample selection in Section~\\ref{sec:data} and the period detection method in Section~\\ref{sec:per_det}. In Section~\\ref{sec:res} we present our results, and examine the mass-period and temperature-period distributions in Section~\\ref{sec:per_mas_teff}, and the amplitude of periodic variability in Section~\\ref{sec:amp_per}. In Section~\\ref{sec:gyro} we compare our results to empirical gyrochronology relations, and in Section~\\ref{sec:comp} we compare our results to those of \\cite{rei+13} and \\cite{nie+13}. We summarize our results in Section~\\ref{sec:disc} and provide full details of the ACF automation method in the Appendix. ",
        "conclusions": "\\label{sec:disc}  The present catalog, of more than 30,000 rotation periods and photometric amplitudes, was derived using a systematic automated search from a well defined sample of more than 133,000 main-sequence stars observed by \\kepler, all of which have stellar temperature estimates. This sample can help us study the mass-age-period relation, and better understand its implication on the amplitude of photometric variability.  One clear feature of the sample is the fraction of rotation period detections as a function of temperature, which changes dramatically, from $\\sim 80$\\% for the cool stars to $\\sim 20$\\% for the hot stars of the sample. It seems that the photometric amplitude of the stellar rotation signature for the majority of the hot stars in the sample is smaller than our detection limit. This is supported by Figure~\\ref{fig:teff_amp}, which shows that the detected amplitudes for the hot stars have a range that goes down to $\\sim 300$\\,ppm, below which we do not detect periodicity. For the cool stars, on the other hand,  the lower limit of the amplitudes is $\\sim 2000$\\,ppm, well above our detection threshold. We might suggest that the incompleteness of our detections for the cool stars is therefore related to the orientation angle of the stellar axis of rotation relative to our line of sight, which prevents us from detecting the rotation of all cool stars in the sample.  Obviously, the photometric amplitude also carries information on the stellar surface inhomogeneity, presumably associated with the stellar magnetic field. Small amplitudes of variability can either result from smaller star spots or from a more even coverage of the stellar surface by spots. Therefore, our sample suggests that the typical inhomogeneity decreases as a function of stellar surface temperature. Further, Figure~\\ref{fig:amp_per} shows that even for a given stellar temperature, the photometric amplitude is a decreasing function of the rotation period. This implies that the stellar inhomogeneity, probably together with the magnetic field, is decreasing as the star is aging, if we adopt the stellar rotation period as an age indicator.  The catalog shows that the typical rotation period increases with decreasing stellar mass and temperature. A comparison to empirical gyrochronology models \\citep{bar07, mam+08, mei+09} shows the upper envelope of points is broadly consistent with an age of 4.5\\,Gyrs, but that the models appear to under-predict the stellar ages, with a significant proportion falling below the 1\\,Gyr rotational isochrone. The position of the Sun close to the upper envelope is interesting, since this implies a lack of stars older than the Sun. More slowly rotating stars should still appear on the main sequence and would not have been removed by our exclusion of evolved stars.  The `dip' feature in the upper envelope around $\\sim~0.5\\,M_{\\odot}$ ($\\sim 4000$\\,K), mentioned in \\cite{mcq+13a} is still visible in the full sample. This mass corresponds to the transition point in young clusters, where lower mass stars spin up considerably more than stars above this mass \\citep{hen+12}. This feature is not accounted for by any current gyrochronology models.  Thorough testing of AutoACF with synthetic LCs containing a range of periods and amplitudes is still required to understand the effect of observational biases on the upper envelope, since low-amplitude, long-period signals are most strongly affected by systematics and quarter joining.  A series of period histograms across varying $T_{\\rm eff}$ bins, shown in Figure~\\ref{fig:hists}, displays some of the finer details of the period distribution, including a significant bimodality in the low-mass stars, which becomes less apparent towards higher masses. This bimodality, forming two sequences in the mass- or $T_{\\rm eff}$-period distribution, cannot be explained by any observational or measurement biases, and thus requires a physical explanation. One hypotheses presented in \\cite{mcq+13a} is that this is a possible hint of two epochs of star formation in the solar neighborhood. This effect may become less apparent towards higher masses, since these brighter stars are seen to greater distances, hence probing a larger region of space.  An alternative explanation, also presented in \\cite{mcq+13a}, is that the dearth of stars between the two sequences represents a mass-dependent, rapid transition phase in the stellar spin-down law, where few stars are observed. In this scenario the star formation rate is unimodal but the spin-down occurs as two slow stages, with a rapid stage in between, producing a bimodal distribution of periods.  Both these hypotheses require further follow-up observations and detailed modeling \\cite[e.g.,][]{gal+13, eps+14}, which will be aided by the larger sample size and wider $T_{\\rm eff}$ coverage of the new data."
    },
    "1402/1402.5970_arXiv.txt": {
        "abstract": "{Atmospheric dispersion and field differential refraction impose severe constraints on widefield, multiobject spectroscopic (MOS) observations, where the two joint effects cannot be continuously corrected. Flux reduction and spectral distortions must then be minimised by a careful planning of the observations -- which is especially true for instruments that use slits instead of fibres. This is the case of VIMOS at the VLT, where MOS observations have been restricted, since the start of operations, to a narrow two-hour range from the meridian to minimise slit losses -- the so-called two-hour angle rule. } {We revisit in detail the impact of atmospheric effects on the quality of VIMOS-MOS spectra with the aim of enhancing the instrument's overall efficiency, and improving the scheduling of observations.} {We model slit losses across the entire VIMOS field-of-view as a function of target declination. We explore two different slit orientations at the meridian: along the parallactic angle (North-South), and perpendicular to it (East-West).} {We show that, for fields culminating at zenith distances larger than 20 deg, slit losses are minimised with slits oriented along the parallactic angle at the meridian. The two-hour angle rule holds for these observations using North-South orientations. Conversely, for fields with zenith angles smaller than 20 deg at culmination, losses are minimised with slits oriented perpendicular to the parallactic angle at the meridian. MOS observations can be effectively extended to plus/minus three hours from the meridian in these cases. In general, night-long observations of a single field will benefit from using the East-West orientation. All-sky or service mode observations, however, require a more elaborate planning that depends on the target declination, and the hour angle of the observations. } {We establish general rules for the alignment of slits in MOS observations that will increase target observability, enhance the efficiency of operations, and speed up the completion of programmes -- a particularly relevant aspect for the forthcoming spectroscopic public surveys with VIMOS. Additionally, we briefly address the (non-negligible) impact of field differential refraction on future widefield MOS surveys.} ",
        "introduction": "\\begin{table}[!h] \\caption{Characteristics of operational and planned/proposed multiobject optical spectrographs with FOVs larger than 100 arcmin$^2$.} \\begin{tabular}{llccccc} \\hline\\hline Telescope & Instrument  & o/p\\,\\tablefootmark{a}  & Aperture & s/f\\,\\tablefootmark{b} &  FOV & N$_{max}$\\\\ &   &  & [m] &  &  [deg$^2$] &  \\\\ \\hline SDSS\t&    \tSDSS     \t\t&   o\t  &\t2.5\t  &     f \t&\t7.07\t\t&\t1000\\\\ WIYN\t&   \tHydraW\t\t&   o   &\t3.5\t  &     f\t&\t0.78\t\t&\t100\\\\ Blanco\t&   \tHydraB\t\t&   o   &\t3.9\t  &     f\t&\t0.35\t\t&\t138\\\\ AAT\t    \t&\t2dF\t\t \t&   o   &\t3.9\t  &     f\t&\t3.14\t\t&\t392\\\\ LAMOST\t&   \tLAMOST\t\t&   o   &\t4.0\t  &     f\t&\t19.6\t\t&\t4000\\\\ WHT\t    \t&\tWYFFOS\t\t&   o   &\t4.2\t  &     f\t&\t0.35\t\t&\t150\\\\ Magellan   &\tIMACS\t\t&   o   &\t6.5\t  &     s\t&\t0.2\t\t&\t600\\\\ MMT\t    \t&\tHectospec\t\t&   o   &\t6.5\t  &     f\t&\t0.78\t\t&\t300\\\\ VLT\t    \t&\tVIMOS\t\t&   o   &\t8.2\t  &     s\t&\t0.08\t\t&\t600\\\\ VLT\t    \t&\tFLAMES\t\t&   o   &\t8.2\t  &     f\t&\t0.14\t\t&\t132\\\\ Mayall\t&   \tDESI\t\t\t&   p   &\t3.9\t  &     f\t&\t7.07\t\t&\t5000\\\\ VISTA\t&   \t4MOST\t\t&   p   &\t4.1\t  &     f\t&\t7.07\t\t&\t3000\\\\ WHT\t    \t&\tWEAVE\t\t&   p   &\t4.2\t  &     f\t&\t3.14\t\t&\t1000\\\\ VLT\t    \t&\tMOONS\t\t&   p   &\t8.2\t  &     f\t&\t0.14\t\t&\t500\\\\ Subaru\t&   \tPFS\t\t\t&   p   &\t8.3\t  &     f\t&\t1.33\t\t&\t2400\\\\ ngCFHT\t&   \tngCFHT\t\t&   p   &\t10.0\t  &     f\t&\t1.5\t\t&\t4000\\\\ GMT\t    \t&\tGMACS\t\t&   p   &\t21.9\t  &     s\t&\t0.05\t&\t400\\\\ \\hline \\end{tabular} \\label{table:facilities} \\tablefootmark{a}{(o)perational/(p)lanned/(p)roposed}\\\\ \\tablefootmark{b}{(s)lits/(f)ibres}\\\\ \\end{table}  The effects of atmospheric dispersion on spectrophotometric observations were first tackled in a seminal paper by \\citet{Filippenko1982}, and have since been addressed by many other authors \\citep{Cohen1988,Donnelly1989,Cuby1994,Szokoly2005}. Two different components contribute to slit losses: a chromatic dispersion caused by the wavelength variation of the index of refraction of air; and an achromatic differential refraction due to airmass variations across the field-of-view (FOV). As pointed out by \\citet{Cuby1998}, the chromatic effect is almost constant for a given field, and can thus be counterbalanced with an atmospheric dispersion compensator (ADC). On the other hand, field differential refraction cannot be continuously corrected, and this is especially problematic for optical instruments with large FOVs. In the case of multiobject spectroscopic (MOS) observations the two joint effects cannot be compensated, so that aperture losses must be minimised by a careful planning of the observations \\citep[e.g.,][]{Cuby1998,Szokoly2005}. This may range from frequent reconfiguration of the fibres, to imposing limited observability windows. The latter is actually the only alternative for instruments that use slits instead of fibres, even more so because field rotation prevents the alignment of all slits along the parallactic angle.  VIMOS\\,\\footnote{http://www.eso.org/sci/facilities/paranal/instruments/vimos.html} \\citep{LeFevre2003} is a widefield (4x7x8 arcmin$^{2}$) instrument with imaging, integral field, and MOS capabilities mounted at the Nasmyth B focus of VLT UT3. The instrument operates in the optical wavelength range (360-1000 nm), and is equipped with six sets of grisms, six sets of broad-band filters, plus three additional filter sets specifically designed to be used in combination with the grisms to block the second order spectra. The unique combination of instrument FOV, large collecting power, and very high multiplexing (up to $\\sim$\\,600 targets in the low resolution modes; see Table\\,\\ref{table:facilities} and Fig.\\,\\ref{fig:vimos}) has made VIMOS a particularly efficient instrument for large spectroscopic surveys of cosmological fields \\citep[e.g.,][]{LeFevre2005,Lilly2007,Popesso2009,Guzzo2013}.  VIMOS-MOS observations are carried out using multislit masks, which provide very accurate sky subtraction and high instrumental throughput.  However, the lack of ADCs makes atmospheric dispersion and field differential refraction factors that need to be taken into account. The atmospheric effects in instruments like VIMOS were studied by \\citet{Cuby1998}. They show that i) atmospheric dispersion dominates at shorter wavelengths, while differential refraction is not negligible at the red end of the visible spectrum; ii) in the former case, image drifts occur along the meridian; iii) differential refraction can be almost neglected for zenith distances smaller than 25 deg for exposure times up to two hours from the meridian (HA = 0\\,h). In view of these results they recommended that, in order to minimise slit losses, the slits be positioned along the dispersion direction at mid-exposure, and observations be limited to a narrow two-hour range from meridian crossing -- the so-called two-hour angle rule. This guarantees that losses remain below 20 per cent for zenith angles $<50$ deg at culmination. These rather limiting guidelines had always been in place for all MOS observations since the start of operations in 2003, and in practice translate into mandatory airmass constraint limits for the VIMOS-MOS observing blocks (OBs).  During the last few years the instrument performance has been significantly enhanced (see \\citealt{Hammersley2010,Hammersley2013}) by changing the detectors to red-sensitive, low-fringing CCDs; replacing the HR-blue grism set with higher throughput VPH grisms; introducing an active flexure compensation system; redesigning the focusing mechanism and mask cabinet; and introducing a new pre-image-less MOS mode \\citep{Bristow2012}. All these improvements have made VIMOS a much more stable instrument, and have extended its lifetime in order to prepare it for the start of the spectroscopic public surveys for which ESO has recently issued a call. Further work to improve the operational efficiency of the instrument includes the present study, which has as main goal to revisit the need for restricted observability of targets only within plus/minus two hours from the meridian in the MOS mode. Increasing the observability of targets in the MOS mode provides more flexibility to operations, because the number of masks that can be loaded in the instrument before the beginning of each night is limited. As we previously noted, VIMOS is very often used for deep observations of cosmological fields, where very long integrations are taken for the same field. By increasing the target visibility (relaxing the two-hour angle rule), the programmes can be completed faster.  \\begin{figure} \\centering \\includegraphics[width=.5\\textwidth]{survey_spectrographs_eff_vimos} \\caption{Survey efficiency (telescope area $\\times$ maximum number of allocated targets) and sky mapping efficiency (telescope area $\\times$ instrument field-of-view, or \\'etendue) of all operational (black), and planned/proposed (grey) multiobject optical spectrographs with FOVs larger than 100 arcmin$^2$. All figures are normalised to those of the SDSS. Note the extraordinary survey efficiency of VIMOS, only comparable to that of next generation instruments.} \\label{fig:vimos} \\end{figure}  In this paper we revisit in detail the impact of atmospheric dispersion on the quality -- in terms of slit losses and spectrophotometric distortions -- of VIMOS observations. % Our aim is to establish general rules for the alignment of slits in MOS observations that enhance the efficiency of operations, not only for VIMOS, but for other operational and future multislit spectrographs as well. We note that the parameter space of this problem is large. Irrespective of image quality and atmospheric conditions, slit losses depend on slit size, orientation, and position within the FOV, observed wavelength range, target declination, total exposure time, and hour angle (HA) of the observations. We build upon the previous work by \\citet{Cuby1998}, but include specific calculations for all the grisms currently available in VIMOS, and present a more detailed analysis of the optimal slit orientation as a function of target declination.  This paper is organised as follows. In Section\\,\\ref{sect:model} we describe the model with which we compute slit losses in VIMOS. Section\\,\\ref{sect:slitlosses} presents the main results of our analysis, while in Section\\,\\ref{sect:fdr} we briefly address the specific contribution of field differential refraction, and its importance for widefield MOS observations. Finally, in Section\\,\\ref{sect:conclusions} we discuss our main results, and summarise our findings. ",
        "conclusions": "\\label{sect:conclusions}  We have shown that the optimal slit alignment in widefield MOS observations is a complex function of target declination and hour angle (Fig.\\,\\ref{fig:nsew}). The traditional orientation of slits along the parallactic angle at the meridian provides very stable results for observations within two hours from the meridian. A setup where slits are oriented \\emph{perpendicular} to the parallactic angle at the meridian is however preferred in the case of fields culminating at low zenith distances, or generally when observations extend to $|$HA$| >$ 2\\,h. This result can readily be understood from the following qualitative line of arguments. As discussed by \\citet{Szokoly2005}, slit losses depend on the interplay between two effects. The first one is independent of the slit orientation, as both atmospheric dispersion and differential refraction are larger at increasing airmasses. On the other hand, the angle between the slit and the parallactic angle is constantly changing. When slits are oriented along the parallactic angle at the meridian (N-S), the angle between the slits and the parallactic angle tends to increase, and the two effects amplify each other. Conversely, when slits are oriented perpendicular to the parallactic angle at the meridian (E-W) --and dispersion is always minimal--, the slits get systematically closer to the parallactic angle at higher airmasses, so that a large fraction of the dispersed light falls back into the slits. This is what creates the opposing trends with declination and hour angle depicted in Figs.\\,\\ref{fig:hrblue}-\\ref{fig:lrred}: the N-S orientation depends strongly on HA, but not so much on $\\delta$, while the strong dependence of  the E-W orientation on $\\delta$ is almost invariant to changes on HA.  \\begin{figure}% \\includegraphics[width=.5\\textwidth]{vimos_airmass} \\caption{VIMOS airmass constraints for MOS observing blocks. The two different shaded areas correspond to the limits for fields that can be observed with slits having N-S orientations at meridian (grey), or E-W orientations (purple). The lower curve in each case corresponds to the physical limit for targets of that declination.} \\label{fig:airmass} \\end{figure}  This has important implications for the scheduling of observations. Night-long observations of a single field (visitor mode) will generally benefit from using the E-W orientation for slits at the meridian. All-sky or service mode observations, however, require a more elaborate planning that depends on the target and the hour angle of the observations. This is a particularly relevant aspect for the forthcoming spectroscopic public surveys with VIMOS, where observations will be carried out in visitor mode. We conclude that the two-hour angle rule, together with the default N-S slit orientation, provide the most stable results for service observations with VIMOS, with slit losses and spectral distortions below 20 per cent -- and almost independent of target declination. This should always be the preferred option for users having targets at $\\delta \\gtrsim -5$ deg or $\\delta \\lesssim -45$ deg. However, for targets within the $-45 \\lesssim \\delta \\lesssim -5$ deg range, the E-W orientation results in comparable or even reduced slit losses. This slit orientation allows for observations to go past the two-hour angle rule, and  be effectively extended up to $|$HA$| = 3$ hours, hence making it a preferred option. This holds for all grisms currently offered in VIMOS, provided the acquisition is done with a filter that closely matches the grism wavelength range.  Figure\\,\\ref{fig:airmass} shows the new airmass constraint limits for MOS OBs. They have been significantly relaxed for fields culminating at small zenith distances, thus increasing target observability. This shall enhance the efficiency of operations, and speed up the completion of programmes. Even though the results presented here are specifically tailored to VIMOS, these general recommendations apply to all current and future widefield multislit spectrographs (e.g., Magellan/IMACS, \\citealt{dressler2011}; GMT/GMACS, \\citealt{depoy2012}). They are already in place since September 2013 for VIMOS-MOS observations \\citep{rsj2013}."
    },
    "1402/1402.0849_arXiv.txt": {
        "abstract": "We report on the discovery of SN~2014J in the nearby galaxy M~82. Given its proximity, it offers the best opportunity to date to study a thermonuclear supernova over a wide range of the electromagnetic spectrum. Optical, near-IR and mid-IR observations on the rising lightcurve, orchestrated by the intermediate Palomar Transient Factory (iPTF), show that SN~2014J is a spectroscopically normal Type Ia supernova, albeit exhibiting high-velocity features in its spectrum and heavily reddened by dust in the host galaxy.  Our earliest detections start just hours after the fitted time of explosion.  We use high-resolution optical spectroscopy to analyze the dense intervening material and do not detect any evolution in the resolved absorption features during the lightcurve rise. Similarly to other highly reddened Type Ia supernovae, a low value of total-to-selective extinction, $R_V \\lsim 2$, provides the best match to our observations.  We also study pre-explosion optical and near-IR images from HST with special emphasis on the sources nearest to the SN location. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are among the most luminous transient events at optical wavelengths and extremely valuable tools to measure cosmological distances, see \\citet{2011ARNPS..61..251G} for a recent review. Yet, SNe~Ia close enough to allow for detailed scrutiny of their physical properties are very rare, especially in a galaxy like M~82, the host of several recent core-collapse supernovae \\citep{2013MNRAS.431.2050M,2013MNRAS.431.1107G}. At an estimated distance to M~82 of 3.5 Mpc \\citep{2009ApJS..183...67D}, SN~2014J is the closest identified Type Ia SN in several decades, possibly rivaled by SN~1972E in NGC~5253 \\citep{1973A&A....28..295A} and SN~1986G in NGC~5128 \\citep{1987PASP...99..592P}.  Thus, SN~2014J is exceptionally well-suited for follow-up observations in a wide range of wavelengths, from radio to gamma-rays. These have the potential to yield transformational new clues into the progenitor systems of SNe~Ia, as well as the detailed properties of dust along the line of sight, key astrophysical unknowns for the study of the accelerated expansion of the universe.  There is strong evidence that SNe~Ia arise from thermonuclear explosions of carbon-oxygen white dwarfs (WD) in binary systems \\citep{2011Natur.480..344N,2012ApJ...744L..17B}. However, the nature of the second star remains unclear. For a long time, the preferred scenario was the single degenerate (SD) model \\citep{1973ApJ...186.1007W}, where a WD accrets mass from a hydrogen or helium rich donor star, thus becoming unstable while approaching the Chandrasekhar mass. The double-degenerate (DD) model involving the merger of two WDs \\citep{1981NInfo..49....3T,1984ApJS...54..335I,1984ApJ...277..355W} has gained considerable observational support in recent years, see e.g., \\citet{2012NewAR..56..122W}.  In this work, we search for potential signatures of a SD progenitor system, such as variable Na~D lines, precursor nova eruptions, features in the early lightcurve, radio emission, or a coincident source in pre-explosion in HST images. ",
        "conclusions": "The discovery of SN~2014J presents us with a unique opportunity to explore the physics of Type Ia SNe and the line-of-sight effects due to intervening matter. Further understanding in these areas is of utmost importance for the use of SNe~Ia in cosmology.  The early data from P48, starting as early as only hours from the explosion, and the multi-wavelength follow-up by the iPTF team covers an important range of the available windows in the electromagnetic spectrum.  Just as the lightcurve reaches its maximum, we have learned that the SN has suffered non-standard extinction. We have searched for, but not detected, any time variation in our high-resolution spectra of the Na~I~D doublet. Similarly, we do not detect any pre-explosion activity in the $\\sim$1500 days of P48 monitoring. In a study of pre-explosion HST images in the near-IR, the nearest resolved source is found $0.2{\\arcsec}$ away from the SN location. The source brightness and offset from the SN makes it unlikely as a donor star in a single-degenerate scenario.  Further, we make a first study of the spectral features of SN~2014J and find that it exhibits high-velocity features from intermediate mass material but lacks C and O often seen in very early spectra. Otherwise, it is a very similar to several well-studied normal SNe Ia."
    },
    "1402/1402.0880_arXiv.txt": {
        "abstract": "Atmospheric neutrinos are produced in interactions of cosmic rays with Earth's atmosphere. At very high energy, the contribution from semi-leptonic decays of charmed hadrons, known as the prompt neutrino flux, dominates over the conventional flux from pion and kaon decays. This is due to the very short lifetime of the charmed hadrons, which therefore do not lose energy before they decay. The calculation of this process is difficult because the Bjorken-$x$ at which the parton distribution functions are evaluated is very small. This is a region where QCD is not well understood, and large logarithms must be resummed. Available parton distribution functions are not known at such small $x$ and extrapolations must be made. Theoretically, the fast rise of the structure functions for small $x$ ultimately leads to parton saturation.  This contribution describes the ``ERS'' \\cite{ers} calculation of the prompt neutrino flux, which includes parton saturation effects in the QCD production cross section of charm quarks. The ERS flux calculation is used by e.g.\\ the IceCube collaboration as a standard benchmark background. We are now updating this calculation to take into account the recent LHC data on the charm cross section, as well as recent theoretical developments in QCD. Some of the issues involved in this calculation are described. ",
        "introduction": "When cosmic rays collide with nuclei in the air in Earth's atmosphere, hadrons are produced in inelastic collisions. These hadrons interact, lose energy, and finally decay, leading to a cascade of particles in the atmosphere that will finally hit the ground. Semileptonic decays of hadrons in the atmosphere generate a flux of neutrinos known as the atmospheric neutrino flux. This is an irreducible background for neutrino observatories such as Super-Kamiokande, IceCube, Antares or the planned KM3NeT. The fluxes at lower energies are pretty well understood, and come mainly from decays of long-lived charged pions and kaons, which are produced in essentially every inelastic collision. This ``conventional'' component of the flux falls steeply with increasing energy, both due to the spectral shape of the incoming cosmic ray flux, and due to the energy loss experienced by the mesons before they decay. Charged pions, for example, have a proper decay length of 8 meters, and with time dilation, their interaction lengths are much smaller than their decay lengths so that they have plenty of time to lose energy. The resulting neutrino energies are therefore downgraded compared to the incoming cosmic ray flux. Theoretical predictions by Honda and collaborators~\\cite{honda,gaisserhonda} of the conventional flux agree very well with measurements up to energies of roughly $10^{5}$~GeV~\\cite{Schukraft:2013ya}.  At even higher energies, production of hadrons containing charm and bottom quarks leads to another component of the flux known as the ``prompt flux.'' These hadrons are produced much more rarely, but since they decay promptly without losing much energy, the resulting flux is harder and has more or less the same energy dependence as the incoming cosmic ray flux. At energies $\\sim 10^{5}-10^{6}$~GeV, this flux is believed to start to dominate the conventional flux, as shown in Fig.~\\ref{fig:fluxes}. \\footnote{The importance of charm production in atmospheric cascades has been known a long time. We have also recently included charm production in the calculation of neutrino fluxes from astrophysical sources \\protect\\cite{Enberg:2008jm}.} Coincidentally, this is roughly the same energy that is starting to be probed by the IceCube experiment. The prompt flux is much more poorly known than the conventional flux, due to the theoretical difficulties involved in calculating charm quark production at very high energies. The calculation is sensitive to very small Bjorken-$x$ and very forward rapidities, and this kinematic region is not probed by present collider or fixed target data. One of the most recent predictions of the prompt flux is the ``ERS'' flux of Ref.~\\cite{ers}, which was computed before the start of the LHC. This prediction is used by the IceCube experiment as a benchmark for the prompt flux.  The IceCube experiment has recently observed two very high energy neutrino events at about 1~PeV~\\cite{Aartsen:2013bka}, and later 26 more events at slightly lower energies~\\cite{Aartsen:2013jdh}, all at and above the energies where the prompt flux becomes important. If these events are to be interpreted as (partly) due to a flux of astrophysical neutrinos, it must be clear that they can not be explained by the background coming from atmospheric neutrinos alone. The theoretical uncertainty in the prompt flux would probably not allow assigning the IceCube observations to atmospheric neutrinos---one would need a prompt flux larger than the ERS flux by a factor 15 in order to get a 10\\% probability to observe two atmospheric events at 1~PeV~\\cite{Aartsen:2013bka}.  However, the significance of their observation depends on the normalization of the atmospheric neutrino flux; the significance of the first two events was $2.8\\sigma$, but with an increase in the prompt flux by a factor 3.8, the significance is reduced to $2.3\\sigma$. Similarly, the significance of the 28 events to not be completely of atmospheric origin is $4.1\\sigma$, which is reduced to $3.6\\sigma$ if the prompt flux background is increased by a factor 3.8. (The factor 3.8 is the level of the measured upper limit on the prompt flux.) Such an increase in the prediction if the various theoretical uncertainties would be better understood is not completely unlikely. On the other hand, if the prediction would become smaller, the significance would increase. In any case, to understand these events it is necessary to understand the background.  The LHC experiments have now measured the charm production cross section at several energies~\\cite{LHCcharm}, providing some constraints on the input to the flux calculation. There have also been theoretical developments in QCD after Ref.~\\cite{ers}, that can be used to constrain the calculation~\\cite{NLL,Albacete:2010sy,Kutak:2012qk}. It is therefore time to improve the calculation to get a better handle on the prediction.  In this contribution I will discuss some of the ingredients in the calculation, and in particular I will discuss why there are large theoretical uncertainties involved. I will necessarily have to skip a lot of details of the calculation, but everything that is not discussed here can be found in our paper \\cite{ers} or in the earlier calculations of the prompt flux \\cite{Lipari:1993hd,Gondolo:1995fq,Bugaev:1998bi,Pasquali:1998ji,Martin:2003us}. See also the book by Gaisser~\\cite{GaisserBook} for details on the description of the cascade in the atmosphere. ",
        "conclusions": ""
    },
    "1402/1402.6696.txt": {
        "abstract": "We study a class of simplified dark matter models in which dark matter couples directly with a mediator and a charged lepton. This class of Lepton Portal dark matter models has very rich phenomenology: it has loop generated dark matter electromagnetic moments that generate a direct detection signal; it contributes to indirect detection in the cosmic positron flux via dark matter annihilation; it provides a signature of the same-flavor, opposite-sign dilepton plus missing transverse energy at colliders.  We determine the current experimental constraints on the model parameter space for Dirac fermion, Majorana fermion and complex scalar dark matter cases of the Lepton Portal framework. We also perform a collider study for the 14 TeV LHC reach with 100 inverse femtobarns for dark matter parameter space. For the complex scalar dark matter case, the LHC provides a very stringent constraint and its reach can be interpreted as corresponding to a limit as strong as two tenths of a zeptobarn on the dark matter-nucleon scattering cross section for dark matter masses up to 500 GeV. We also demonstrate that one can improve the current collider searches by using a Breit-Wigner like formula to fit the dilepton MT2 tail of the dominant diboson background. ",
        "introduction": "\\label{sec:intro} %__________________________________________________% The search for thermal relic Weakly Interacting Massive Particle (WIMP) dark matter has a long history, particularly within models of weak-scale supersymmetry (SUSY)~\\cite{Haber:1984rc,Jungman:1995df}. Such models can furnish both signatures of new physics at the TeV scale and a viable candidate for dark matter.  Collider, direct detection and indirect detection searches for Minimal Supersymmetric Standard Model (MSSM) dark matter particles have limited the vanilla parameter space~\\cite{CahillRowley:2012kx,Cahill-Rowley:2013dpa} and weaken the strong tie between WIMP dark matter and the SUSY framework.  Outside of weak-scale SUSY models, there is no specific reason for dark matter to have mass near 100 GeV. Fortunately, the ``WIMP miracle'' provides guidance for the plausible region of dark matter mass and interaction strength~\\cite{Kolb:1990vq}.  Since even the discovery of dark matter from multiple experimental probes is unlikely to immediately tell us the underlying framework, in this paper we concentrate on a class of simplified dark matter models, which serves as a phenomenological bridge between experiments and a deep underlying theory.  There have been a number of recent studies of simplified dark matter models with the emphasis on the complimentarity from different experimental searches~\\cite{Chang:2013oia,An:2013xka,Bai:2013iqa,DiFranzo:2013vra,Buchmueller:2013dya,Cheung:2013dua,Papucci:2014iwa,Simone:2014}. Most of those studies have concentrated on dark matter interactions with the quarks of the Standard Model (SM), which leads to a new framework for interpretation of LHC and direct detection searches in terms of dark matter properties. For instance, in Ref.~\\cite{Bai:2013iqa} the signature of two jets plus missing transverse energy has been studied within the context of Quark Portal dark matter models, which is a class of simplified models in which dark matter particles and mediators interact with a single quark. In this paper, following our previous study in Ref.~\\cite{Bai:2013iqa}, we concentrate on the lepton sector and study a class of Lepton Portal dark matter models.  In these models, there are two new particles in the dark matter sector with the lightest one being the dark matter candidate, which must be a singlet under electromagnetism and color.  The other particle plays the role of mediator and connects the dark matter particle to the leptons. Obviously, to conserve SM gauge symmetry, the mediator particle should be charged under the electroweak symmetry.  For the dark matter interactions to be renormalizable, the mediator must have the same quantum numbers as the left-handed lepton weak doublet or the right-handed charged leptons. In our study, we consider only the latter case for simplicity.  Compared to Quark Portal dark matter models, Lepton Portal dark matter models have totally different phenomenology at the three frontiers of the search for WIMP dark matter.  For direct detection, unlike the Quark Portal case, dark matter particles do not directly couple to target nuclei at tree level.  At one loop, the dark matter can couple to the photon through various electromagnetic moments, which generates the dominant interaction with the target nucleus.  The latest LUX results from Ref.~\\cite{Akerib:2013tjd} can constrain a large portion of parameter space for Dirac fermion or complex scalar dark matter. For indirect detection, dark matter particle annihilation can generate electrons or positrons with a harder spectrum than the Quark Portal case. Hence, the electron and positron flux measurement from AMS-02 in Refs.~\\cite{Aguilar:2013qda,AMS-electron-positron} becomes relevant for the Lepton Portal models. At colliders, the Quark Portal models have a larger signal production but also a larger QCD background. In the Lepton Portal models, the dark matter mediator particles can be pair produced via off-shell photons or $Z$ bosons. The corresponding collider signature is two same-flavor charged leptons plus missing transverse energy. Because both ATLAS and CMS collaborations at the LHC can make very good measurements of charged lepton momenta, the signature of dilepton plus missing transverse energy could serve as the \\emph{discovery channel} for dark matter particles. Therefore, we pay more attention to understanding and optimizing the key kinematic variables and work out the sensitivity at the 14 TeV LHC.  Colliders can cover the light dark matter mass region beyond the direct and indirect detection sensitivity. This is simply due to different kinematics for different probes. For the three categories of dark matter particles: Majorana fermion, Dirac fermion and complex scalar, we have found that the 14 TeV LHC has a much better reach than the direct detection experiments for the Majorana fermion and complex scalar cases.  For the Majorana case, the dark matter scattering cross section is suppressed by the dark matter velocity and predicts a very small rate for direct detection experiments. For the complex scalar case, the dark matter fermion partner has a large production cross section at the LHC and a high discovery probability at the LHC.  Our paper is organized as follows. In Section \\ref{sec:lepton}, we introduce the Lepton Portal class of simplified models.  We determine the allowed parameter space for dark matter to be a thermal relic in Section~\\ref{sec:relic-abundance}. The direct detection will be covered in Section~\\ref{sec:direct-detection}, where we perform loop-level calculations to determine the dark matter elastic scattering cross section. In Section~\\ref{sec:indirect-detection}, we work out constraints on model parameter space from the AMS-02 positron and electron flux measurement. We then perform a collider study for the sensitivity at the 14 TeV LHC with 100 fb$^{-1}$ and present summary plots in Section~\\ref{sec:collider}.  We conclude in Section~\\ref{sec:conclusion}.   %__________________________________________________% ",
        "conclusions": "\\label{sec:conclusion} %__________________________________________________% We want to first emphasize the importance of colliders for discovering or excluding the Lepton Portal dark matter.  The signature with the same-flavor and opposite-sign dilepton plus missing energy is a pretty clean one. The $M_{T2}$ cut can be imposed to make almost background free. As a result, the discovery reach is purely determined by the signal production cross section times the acceptance. For a large mass splitting between dark matter and its partner, the signal acceptance is large, so the discovery reach is limited by the signal cross section. From Fig.~\\ref{fig:feyn-diagram}, one can see a large increase of the mediator production cross sections from 8 TeV to 14 TeV and a discovery of dark matter signals at the LHC may happen in the near future.  In our analysis, we have considered both the electron and muon cases and neglected the tau lepton case. We anticipate a slightly weaker limit from the LHC because of the tau-tagging and mis-tagging efficiencies.  Another parameter region that we have ignored is the co-annihilation region.  The collider searches become less sensitive because the leptons from the mediator decays are either too soft to pass the basic cuts or generate insufficient $M_{T2}$ and would be buried in the SM backgrounds.  In the extremely degenerate region, one could include an additional jet, photon, $W$ and $Z$ gauge bosons from initial state radiation to gain sensitivity.  In summary, we have studied Lepton Portal dark matter for three cases: Majorana fermion, Dirac fermion and complex scalar dark matter. For direct detection, the majorana fermion case has a very small predicted event rate because of the leading operator of the dark matter coupling to photon has an additional velocity suppression. On the other hand, the direct detection signals for the Dirac fermion and complex scalar cases are not suppressed. In terms of indirect detection, since only the Dirac fermion case has non-zero $s$-wave annihilation, AMS-02 has the best coverage for its model parameter space. At colliders, the LHC has better reaches for the light dark matter mass region than the direct detection experiments.  For the complex scalar case, the 14 TeV LHC with 100 fb$^{-1}$ can cover mediator masses up to 800 GeV and provides a constraint on spin-independent dark matter-nucleon scattering cross section as low as $2\\times10^{-46}$~cm$^2$ for dark matter masses up to 500 GeV and a unit coupling.  \\subsection*"
    },
    "1402/1402.2631.txt": {
        "abstract": "Massive stars can alter physical conditions and properties of their ambient interstellar dust grains via radiative heating and shocks. The \\ion{H}{2} regions in the Large Magellanic Cloud (LMC) offer ideal sites to study the stellar energy feedback effects on dust because stars can be resolved, and the galaxy's nearly face-on orientation allows us to unambiguously associate \\ion{H}{2} regions with their ionizing massive stars. The \\emph{Spitzer Space Telescope} survey of the LMC provides multi-wavelength (3.6--160 $\\mu$m) photometric data of all \\ion{H}{2} regions. To investigate the evolution of dust properties around massive stars, we have analyzed spatially-resolved IR dust emission from two classical \\ion{H}{2} regions (N63 and N180) and two simple superbubbles (N70 and N144) in the LMC. We produce photometric spectral energy distributions (SEDs) of numerous small subregions for each region based on its stellar distributions and nebular morphologies. We use \\texttt{DustEM} dust emission model fits to characterize the dust properties. color--color diagrams and model fits are compared with the radiation field (estimated from photometric and spectroscopic surveys). Strong radial variations of SEDs can be seen throughout the regions, reflecting the available radiative heating. Emission from very small grains drastically increases at locations where the radiation field is the highest, while  polycyclic aromatic hydrocarbons (PAHs) appear to be destroyed. PAH emission is the strongest in the presence of molecular clouds, provided that the radiation field is low. ",
        "introduction": "\u00ca\\label{sec:intro} %=================================================================  Far infrared (FIR) emission of interstellar dust has been commonly used to gauge star formation rates \\citep[e.g.,][]{Cetal10}. \u00caIn a star-forming region, interstellar dust can be heated, sputtered, and shattered by energy feedback from massive stars via ultraviolet (UV) radiation, fast stellar winds, and supernova remnant (SNR) shocks. \u00caIt would thus be prudent to investigate how dust properties are modified in the harsh environment around massive stars. \u00ca  The Large Magellanic Cloud (LMC) provides a nearby, ideal laboratory to study the influence of massive stars on dust properties because of the following advantages over our Galaxy: (1) the LMC has a nearly face-on orientation, mitigating the confusion and extinction along the Galactic plane; (2) the LMC is at a known distance, $\\sim$50 kpc \\citep{fea99}, so stars can be resolved and studied in conjunction with the interstellar gas and dust; and (3) there exist many complementary multi-wavelength observations of stars, gas, and dust in the LMC. For these reasons, the LMC has been targeted by every space-based IR observatory to study dust properties and calibrate dust emission as star formation indicators.  The first report of \\emph{Spitzer Space Telescope} \u00caobservations of an LMC \\ion{H}{2} complex was made by \\citet{Getal04} for LHA\\,120-N206 \\citep[N206 for short; designation from][]{hen56}; its dust emission was qualitatively compared with that of the Orion Nebula. \u00caSubsequently, the entire LMC was surveyed by \\emph{Spitzer} in the Legacy program Surveying the Agents of a Galaxy's Evolution \\citep[SAGE;][]{Meixner06}. \u00ca\u00caUsing the SAGE \u00caand additional data of 32 \\ion{H}{2} complexes in the LMC and Small Magellanic Cloud (SMC), \\citet{Letal10} found that the 70 $\\mu$m fluxes appear to be the best star formation indicator, although the correlation between the 70 $\\mu$m flux and the bolometric IR flux in the 30 Dor giant \\ion{H}{2} region shows a large degree of scatter. \u00caThe anomaly in 30 Dor could be attributed to the intense energy feedback from the large number of massive stars encompassed by this region.  The SAGE data have also been used to analyze the global FIR emission in the LMC and spectral energy distributions (SEDs) of three large regions selected to represent different levels of star-formation activity \\citep{Betal08}. A 70 $\\mu$m excess in the SEDs is found, and this excess increases along the sequence ``Milky Way--LMC--SMC'', suggesting an influence of decreasing metallicity. \u00ca\\citet{Betal08} fit the SEDs with an improved version of the dust emission model by \\citet{Detal90} and suggested that the 70~$\\mu$m excess can be explained by modifying the size distribution of very small grains (VSGs) to have an increased abundance of larger VSGs. Since large VSGs can be produced through erosion of large grains in the diffuse medium, the presence of VSGs suggests effects of energy feedback.  To investigate the effects of massive stars on dust emission, \\citet{Slater11} studied the dust properties of 4 classical and 12 superbubble \\ion{H}{2} regions in the LMC and found little correlation between dust emission properties and the spectral type of the hottest star, the bolometric luminosity of the underlying massive stars, or the evolutionary status of the \\ion{H}{2} regions. \u00caThis result is surprising and is at odds with the anomalous properties of 30 Dor and the abundant presence of large VSGs.  We have examined the \\ion{H}{2} region sample of \\citet{Slater11} and find that their division between classical and superbubble cases may not be consistent with observations. \u00caFor example, DEM\\,L34 \\citep[= N11; DEM designation from][]{DEM76} and DEM\\,L152 (= N44) both contain multiple OB associations with a central superbubble surrounded by compact \\ion{H}{2} regions, but are classified as ``classical'' and ``superbubble'' \\ion{H}{2} regions, respectively. \u00caFurthermore, they have used the integrated dust emission from entire \\ion{H}{2} regions, while it is expected that the influence of stellar energy feedback ought to depend on radial distances from the stars.  To investigate the influence of massive stars on interstellar dust grains with the least ambiguity, we have chosen to study two clearly defined classical \\ion{H}{2} regions, N63 (= DEM\\,L243) and N180 (= DEM\\,L323), and two relatively simple superbubbles, N70 (= DEM\\,L301) and N144 (= DEM\\,L199), whose stellar content is well studied. We have analyzed \\emph{Spitzer} observations of these four regions, measuring flux densities in radially distributed subregions, constructing SEDs, and using dust emission model fits to determine the dust properties. \u00caIn this paper, we describe our methodology, target selection, and data sets used in Section \\ref{sec:hii}, discuss the stellar content and assess the stellar energy feedback in Section \\ref{sec:SCRF}, describe the creation and modeling of SEDs in Section \\ref{sec:SEDs}, examine the results of each \\ion{H}{2} region in Section \\ref{sec:dust_analysis}, discuss the relationship between stellar energy feedback and dust properties in Section \\ref{sec:disc}, and summarize the findings in Section \\ref{sec:sum}.  % \u00caGalliano et al. 2011, astro-ph \u00ca[Herschel data] \u00ca\\\\ % \u00caI find Galliano et al paper not very useful for our intro section.  %================================================================ ",
        "conclusions": "\\label{sec:disc} %=========================================================== We have used model fits to spatially resolved SEDs of two classical and two superbubble \\ion{H}{2} regions to determine the dust composition and properties with respect to the radiation field. In this section we discuss the individual dust components and compare the dust properties in differing \\ion{H}{2} regions and SNRs.  \\subsection{Dust} \\subsubsection{PAHs}   Radiation is needed to excite PAHs; however, PAHs also can be destroyed in the presence of a strong radiation field \\citep[e.g.,][]{Madden06,Lebout07}, primarily due to high energy photons \\citep[e.g.,][]{Aitken85,Voit92}. Moreover, the location of PAH and H$_2$ emission are strongly correlated in cool-PDRs, such as the filaments in the Horsehead Nebula \\citep{Hab05,com07} and $\\rho$ Ophiuchus \\citep{Hab03}. However, the locations become anti-correlated with higher radiation field, as seen in the Orion bar \\citep{Tie93} and in Monoceros R2 \\citep{Ber09}.  The results in this paper are consistent with these studies. We find that PAH emission is enhanced toward locations of molecular clouds except where the radiation is high. We also find that PAH emission is diminished at locations that are optically thin to UV radiation.  Since our regions are relatively young (no known SNRs in N144 and N180, one known in N63, and a history of one or multiple in N70), the radiation field is dominated by massive stars that produce harsh photons. Although high energy photons are the main mechanism for destroying PAHs rather than the radiation field, we will assume that the stellar content of the regions follow the same initial mass function and have similar spectral shape in their radiation field and attempt to provide an estimate of the typical $U$ required for significant PAH destruction. In particular, we focus on N144 and N180 since N63 lacks large-scale CO observations and the dust emission from N70 is thought to come from the surrounding shell. For these two regions, Figure \\ref{fig:Upah} both $Y_{\\rm{PAH}}/Y_{\\rm{DUST}}$ and the color [8.0]--[24.0] versus the radiation field as ascertained by massive stars ($U$). We only show subregions when at least half the region contains a 2$\\sigma$ Magellanic Mopra Assessment \\citep[MAGMA;][]{Wong11} CO detection. Note that higher values of the color [8.0]--[24.0] indicates less PAH emission. For both panels, subregions for N180 indicate that PAHs are less coincident with CO molecular gas for high $U$, particularly above $U \\approx 30$. N144 subregions have points in the same general vicinity as N180, but an obvious correlation between PAHs and radiation field is not apparent. This could be due to the lack of data (particularly at low values of $U$) or physical differences between N144 and N180.  Our calculations of the radiation field may be highly uncertain since they can be quite dissimilar (sometimes by about an order of magnitude) to the fitted results. However, the fitted results of $U$ are typically below those of the massive stars, so $U \\approx 30$ is a fair estimate of an upper limit of the typical radiation field needed to destroy PAHs in N180. Though the $U$ limit of N180 may not be representative of other molecular clouds, we note that maps in \\citet{Galametz13} also show a reduction at PAHs with similar modeled $U$ values in the LMC classical \\ion{H}{2} region N159, with a drastic depletion in the fraction of PAHs for $U \\gtrsim$ 100.   %Plots in \\citet{Galametz13} indicated a value consistent with ours; the modeled radiation field in the classical \\ion{H}{2} region N159 in the LMC has a drastic depletion in the fraction of PAHs for $U \\gtrsim$ 100.  %This value is roughly consistent with classical \\ion{H}{2} LMC region N159 where the modeled fraction of PAHs is drastically reduced for $U \\gtrsim$ 100 \\citep{Galametz13}.  %In \\citet{Galametz13}, it is seen that in some regions of the LMC, the modeled fraction of PAHs is drastically reduced for $U \\gtrsim$ 100. Our results push this limit a bit lower, though this limit is certainly more dependent on the quantity of high energy photons rather than the value of $U$.   %Dust modeling in \\citet{draine2007} showed that with a fix fraction of dust that are PAHs, for $U \\lesssim 10^3$ and for $\\lambda<20 \\mu$m, the \\emph{normalized} emission for dust grains is independent of $U$. Thus, given a constant PAH mass fraction throughout our subregions, we would expect an increase in PAH emission where the radiation is stronger. However, we certainly see a decrease for increasing $U$, which may suggest that the PAH fraction is not constant across our subregions due to their destruction via radiation.  %As highlighted in \\citet{Ber09}, several studies have analyzed the spatial distribution of PAH emission and rotational H$_2$ lines in PDRs. The location of PAH and H$_2$ are strongly correlated in cool-PDRs, such as the filaments in the Horsehead Nebula \\citep{Hab05,com07} and $\\rho$ Ophiuchus \\citep{Hab03}. However, the locations becomes anti-correlated with higher radiation field, as seen in the Orion bar \\citep{Tie93} and in Monocerous R2 \\citep{Ber09}. Indeed, observations have shown that these PAHs are destroyed due to radiation \\citep[e.g.,][]{Madden06,Lebout07}, particularly by high energy photons \\citep[e.g.,][]{Aitken85,Voit92}.  %\\cite{Ber09} suggested that for a dense PDRs with a high UV field, the rotational H$_2$ lines is caused by collisional excitation and thus highly dependent on gas temperature, causing the locations of PAHs and warm H$_2$ lines to be different.  %\\cite{Ber09} suggested that this spatial discrepancy may be due to a selection bias; for a dense PDRs with a high UV field, the rotational H$_2$ lines is caused by collisional excitation and thus highly dependent on gas temperature, causing the locations of PAHs and warm H$_2$ lines to be different.     %\\citet{Ber09} attributed this fact that in dense PDRs, high UV causes the rotational H$_2$ lines to be caused by collisional excitation  \\subsubsection{VSGs and BGs} \\label{sec:VSGBG} In all of our \\ion{H}{2} regions except N70 (whose dust is uniformly distributed about its shell), $Y_{\\rm VSG}$/$Y_{\\rm BG}$ is highest at locations where the radiation field is the strongest, as evident in Figure \\ref{fig:U_vs_vsg_bg_ratio-eps-converted-to.pdf}. We stress that the VSG population is highly dependent on each SED's 24~$\\mu$m excess and is typically anti-correlated with the BG population.  Grain populations typically change due to both destruction (i.e., dust changes back to gas) or disruption (large grains break into smaller grains). \\citet{Jones96} showed that fast shocks ($>$150 km~s$^{-1}$) likely destroy grains via sputtering, while the velocity requirement for shattering BGs with VSGs depends on the dust species, with velocities ranging from $\\sim$1--3 km~s$^{-1}$. In our regions, dust-dust collisions are probably uncommon due to low densities ($n_{\\rm{H}} \\sim 100$~cm$^{-3}$). However, O-stars are known to have wind-blown bubbles expanding at velocities of 10--15 km~s$^{-1}$ \\citep[with observations of such bubbles in N180,][]{Naze01}, which could lead to grain shattering.  Other possibilities of an increased $Y_{\\rm VSG}$ include dust coagulation and dust erosion. Dust coagulation of the smallest grains into slightly larger VSGs could increase the $24~\\mu$m emission, but the dust in our regions are not in environments conducive to coagulation \\citep[see, e.g.,][]{koh12}. \\citet{Betal08} suggested erosion of BGs into VSGs likely explains the LMC's 70~$\\mu$m excess in neutral/ionized regions. Moreover, dust erosion is thought to be increased during star formation due to increased dust processing \\citep{par11}. The central OB associations of our \\ion{H}{2} regions have a rich history of previous star formation, which could certainly lead to more dust processing/erosion than the areas analyzed at the macro-scale (several 100~pc; \\citealt{Betal08}), leading to the increased 24~$\\mu$m emission. Thus, at scales of $\\sim$10~pc, BGs near OB associations may have gone through substantial erosion that has caused this increase of VSGs.  %Radiation pressure on a VSG depends on the dust absorption efficiency. For grains smaller than the photon wavelengths, this efficiency is very low, making the radiation pressure negligible. However, for the massive stars in our \\ion{H}{2} regions, the dominate radiation wavelengths becomes very similar to the sizes of the VSGs, allowing for the larger VSGs to have dust absorption efficiency near unity \\citep[e.g.,][]{com2011}. A calculation shows that during a VSG's mean free path between collisions, radiation pressure can accelerate VSGs to shattering velocities (a few km/s) without reaching catastrophic speeds \\citep[$\\gtrsim 100 km~s^{-1}$,][]{Jones96}. Moreover, O-stars are shown to have wind-blown bubbles expanding at velocities of 10-15 km~s$^{-1}$ \\citep[including N180,][]{Naze01}. Therefore, at locations where there is a large radiation field in these \\ion{H}{2} regions, it is sensible to see this substantial increase in $Y_{\\rm VSG}$/$Y_{\\rm BG}$.  On the other hand, \\citet{Pilleri2012} saw a definite decrease of ``evaporating\" VSGs with increasing radiation field. The sources they analyzed typically had a higher radiation field ($U \\gtrsim 100$) than our regions, though we still see an increase in VSGs from $U=100$ to $U=1000$. However, their dust model of VSGs follows more closely to the smallest grains of \\citet{Li01} while ours follows \\citet{Detal90}. Since \\citet{Pilleri2012} modeled VSGs that are more similar to PAHs and thus have an emission spectrum at much shorter wavelengths, these VSGs are not comparable to the VSGs modeled in this paper.  %Although radiation can cause the primordial VSGs to be preferentially driven away over BGs in less than 1 Myr, we observe more VSG emission with increasing ISRF. VSGs are potentially replenished by BGs, whose shattering threshold may be achieved via stellar winds from nearby massive young stars.  % What makes the ``shocks'' so important? % High enough dust velocities that can cause shattering? % Can dust be considered an ideal gas? % Shocks in the past? % % (1) radiation pressure drives smaller grains out faster than big grains, % \u00ca\u00ca\u00ca\u00ca making differential speed % (2) fast stellar wind can vaporize the dust grains, if hit % \u00ca\u00ca\u00ca\u00ca I made some crude calculation, and it seems plausible % (3) ionization front drives shocks into the ambient medium  % (1) doesn't make sense to me since more radiation we see more smaller grains % (2) ditto  % Meeting notes from Dec 12, 2012 -- back of the envelope calculation says a 100 angstrom % grain w/ density 2.2 g / cm^3 can be pushed as far as 38 pc (sans collisions) in 1 Myr from a % 10^38 erg / s luminosity. \u00caPossibly useful sentence: Although primordial VSGs can be % radiatively driven from HII regions by massive ionizing stars on the order of 1 Myr, we % observe more VSG emission with increasing ISRF. \u00caVSGs are potentially \u00ca % replenished by BGs, whose shattering threshold may be achieved via stellar winds % from nearby massive young stars.   %\\subsection{Dust properties and massive stars} %\\subsection{Radiation Fields of Regions Analyzed} %\\subsubsection{Spectroscopically Classified Massive Stars and MCPS stars}  %(1) spectroscopically classified massive stars \\\\ %(2) blue stars found from MCPS via CMDs (this description needs to happen in sec 4)\\\\ %Plot these stars and mark them in IR/Ha images.   %Spectral typing can further quantify the stellar radiation field; \\citet{Garmany94} and \\citet{Oey96} provide spectral typing for the most massive stars in these regions. \u00ca\\citet{Oey96} reports 130 typed objects within the LMC, the locations of which (will be plotted soon) are plotted in Figs ...  %\\subsection{SEDs and Stellar Feedback} %To quantify the effects of stellar feedback on the ISM, we pair the overall SED trend and location of each SED region to nearby stellar sources.  \\subsection{Dust Properties of Evolved \\ion{H}{2}~Regions} We find that N70, the most evolved \\ion{H}{2} region, likely has its dust mass located in the shell. However, major differences in dust properties between classical \\ion{H}{2} regions and superbubbles are not obvious with this small sample size. \\citet{Relano13} analyzes the SEDs of 119 \\ion{H}{2} regions in M33, categorizing them in a likely evolutionary sequence (filled, mixed, shell, and clear shell morphologies). They found that each category had similar SED features, though differed slightly between each classification. Specifically, they found that the FIR peak is located at longer wavelengths for older regions (shells and clear shells), indicating that colder dust is likely present. Uniform cold dust temperatures are expected since gas has been pushed far away from the central heating source. This result is consistent with our analysis of N70, which shows a uniform, low fraction of $Y_{\\rm{VSG}}/Y_{\\rm{BG}}$ compared to our less evolved \\ion{H}{2} regions that do not have the same shell geometry.  \\subsection{Comparison to SNR Studies} Changes in dust populations are commonly studied in SNRs due to their dynamic properties. Fast shocks ($\\gtrsim150$ km~s$^{-1}$) allow dust to reach high velocities that may cause their destruction via catastrophic shattering or sputtering \\citep[e.g.,][]{Jones96,Micelotta10}, though VSGs and PAHs can also be created through low-velocity shattering \\citep{Jones96}. While PAHs can survive or even be produced in slow shocks, fast shocks are generally considered to destroy all PAHs in a region except in circumstances where they can be recreated or protected by their environment \\citep[e.g.,][]{Micelotta10,Seok12}. Additionally, harsh photons provide another mechanism of destroying the smallest grains; the radiation field in SNRs can vary drastically, though it is typically about 10-100 times the local ISRF \\citep[e.g.,][]{Andersen11}. With these different mechanisms for creation and destruction of dust, it is important to observe and model a wide range of SNRs with different shock velocities, densities, and radiation fields to test theoretical models. Moreover, dust properties of SNRs can be compared to our \\ion{H}{2} regions to provide a diagnostic of how dust changes in varying radiation fields in the absence of large shock velocities.  %Even at the youngest ages ($<300$ yr), dust is known to be produced in the ejecta of these SNRs \\citep{Rho09}. However, since PAHs were only recently discovered with $Spitzer$ \\citep{Tappe06}, studies of the dust properties have been limited. Since then, \\citet{Andersen11} used \\texttt{DustEM} (also using the \\citealt{Detal90} dust models) to fit BGs, VSGs, and PAHs for 14 Galactic SNR, representing the largest survey to model these properties in SNRs.  The \\emph{Spitzer} IR Spectrograph (IRS) has been used to study a large number of SNRs.  Dust produced in supernova ejecta has been detected in young ($<$300 yr) SNRs \\citep{Rho09}, and PAHs are detected in SNRs \\citep[e.g.,][]{Tappe06}.  Using a large spectral survey of 14 Galactic SNRs, \\citet{Andersen11} used \\texttt{DustEM} (also using the \\citealt{Detal90} dust models) to model the contributions of emission from BGs, VSGs, and PAHs. They found that the ratio of VSGs to BGs is 2--3 times larger in SNRs than the plane of the Galaxy, and they attribute dust shattering as the mechanism for creating the surplus of VSGs. Two of the 14 SNRs had relatively high radiation fields; fits of these regions had relatively low PAHs and virtually no VSGs. \\citet{Andersen11} suggested that these regions have high shock velocities and low densities, which can lead to destruction of VSGs and PAHs through sputtering.  With high radiation, we find that the VSG population can increase and PAHs are generally destroyed; however, fast shocks can destroy both grains. Unlike SNR regions, we find that VSG population increases where there is higher $U$, and there are no known fast shocks in our regions. Indeed, while observations from \\emph{XMM Newton} show X-ray emission from N70, there is no X-ray emission toward N180 and N144, and for N63, X-ray emission is only associated with the known SNR N63A (S. L. Snowden et al. in preparation). It has been suggested that N70 has recently been heated by an interior SNR \\citep{CM90,rod11,zha14}. If shocks are the main mechanism that are heating the grains in N70 rather than the ISRF, our \\texttt{DustEM} modeling would be invalid (since $U$ is our free parameter rather than temperature). Regardless, the shape of the SEDs of each subregion is the same with a dip at 24~$\\mu$m, which agree with our \\texttt{DustEM} fit of a uniform, low abundance of VSG. Therefore, it is most likely that a past history of fast shocks has destroyed the population of VSGs, and no mechanism has been available for creation of new VSGs (due to, e.g., the lack of centrally located gas since it has been blown away). %and that knowledge of shock velocities is crucial when analyzing the dust populations in a region.    %====================================================================="
    },
    "1402/1402.6554_arXiv.txt": {
        "abstract": "{} {In previous studies a very hot plasma component has been diagnosed in solar active regions through the images in three different narrow-band channels of SDO/AIA. This diagnostic from EUV imaging data has also been supported by the matching morphology of the emission in the hot Ca XVII line, as observed with Hinode/EIS. This evidence is debated because of unknown distribution of the emission measure along the line of sight. Here we investigate in detail the thermal distribution of one of such regions using EUV spectroscopic data.} {In an active region observed with SDO/AIA, Hinode/EIS and XRT, we select a subregion with a very hot plasma component and another cooler one for comparison. The average spectrum is extracted for both, and 14 intense lines are selected for analysis, that probe the $5.5 < \\log T < 7$ temperature range uniformly. From these lines the emission measure distributions are reconstructed with the MCMC method. Results are cross-checked with comparison of the two subregions, with a different inversion  method, with the morphology of the images, and with the addition of fluxes measured with from narrow and broad-band imagers.} {We find that, whereas the cool region has a flat and featureless distribution that drops at temperature $\\log T \\geq 6.3$, the distribution of the hot region shows a well-defined peak at $\\log T = 6.6$ and gradually decreasing trends on both sides, thus supporting the very hot nature of the hot component diagnosed with imagers. The other cross-checks are consistent with this result.} {This study provides a completion of the analysis of active region components, and the resulting scenario supports the presence of a minor very hot plasma component in the core, with temperatures $\\log T > 6.6$. } ",
        "introduction": "It is accepted that the energy source that sustains the high temperature of solar corona is in the magnetic field \\citep{2006SoPh..234...41K}. Regarding the way in which this energy is converted to thermal energy, we can broadly distinguish between two scenarios that depend on the timescale of the energy release: one is continuous heating, the other is in the form of discrete, rapid pulses. The latter is consistent with the nanoflare model proposed by  \\cite{1988ApJ...330..474P}. According to this model, the magnetic field tubes are displaced by the photospheric motions, and can approach and interact. When two flux tubes are almost in contact they will form a current sheet, where the field lines can reconnect. The reconnection can release a large quantity of energy in impulsive events called nanoflares. Signatures of these events are difficult to observe for several reasons, one of which is the fine structuring of the magnetic loops, that is hardly resolved with present-day instruments (e.g., \\citealt{2013ApJ...770L...1T}).  One feature to discriminate the heating release is the presence or absence of very hot plasma. In active regions, the mean coronal temperature is typically 2-3 MK. If heat pulses occur, we expect that a small amount of plasma hotter than the average (6-10 MK) will be ever-present.  Recently, a number of studies, mostly based on data from the Hinode and the Solar Dynamics Observatory (SDO) missions, have shown increasing evidence for such small very hot components in active regions (\\citealt{2009ApJ...704L..58R}, \\citealt{2009ApJ...698..756R}, \\citealt{2009ApJ...697...94M}, \\citealt{2009ApJ...693L.131S}, \\citealt{2009ApJ...697.1956K}, \\citealt{2009ApJ...696..760P}, \\citealt{2010AstL...36...44S}, \\citealt{2010A&A...514A..82S}), but the issue is still debated (\\citealt{2012ApJ...754L..40T}, \\citealt{2011ApJ...734...90W}). In SDO images taken with a channel sensitive also to emission of plasma at 6 MK (94 \\AA), cores of active regions contain bright strands \\citep{2011ApJ...736L..16R}, as predicted by models of nanoflaring loops \\citep{2010ApJ...719..576G}. It remains to be proven whether this plasma is really at such high temperatures or not. Spectroscopic observations should help greatly. In this work, we analyse an active region that shows evidence for this very hot component in SDO data, but for which spectroscopic data are also available from the EUV spectrometer EIS on-board the Hinode mission. In a previous work  (\\citealt{2012ApJ...750L..10T}, hereafter Paper I), emission in the CaXVII line, which forms around temperature of 6-8MK, was detected in the hot structures identified with SDO/AIA data. In that work they built a 3 color image, to highlight the presence of a very hot component of emitting plasma inside the active region. The AIA 94~\\AA\\ band  is known to be multi-thermal. It is sensitive to hot plasma, due to the presence of an Fe XVIII line, formed around 6~MK, but is also sensitive to plasma at 1 MK, because of the presence of a Fe X line and  cooler Fe IX and Fe VIII (see, e.g., \\citealt{2011A&A...535A..46D}, \\citealt{2012ApJ...745..111T}, \\citealt{2011ApJ...743...23M}, \\citealt{2011ApJ...740L..52F}, \\citealt{2012A&A...537A..22O}, \\citealt{2012A&A...546A..97D}). Recently, also a Fe XIV line was identified in  \\cite{2012A&A...546A..97D}. This line is normally stronger than the other cool components in active region cores, as shown in \\cite{2013A&A...558A..73D}. It is therefore not simple to assess if the hot emission seen in the AIA 94~\\AA\\ band is really due to Fe XVIII. To clarify this point, \\cite{2012ApJ...750L..10T} compared the AIA 94~\\AA\\ image with the Ca XVII image obtained from the EIS spectrometer. They showed a strong correlation between the hot CaXVII and the emission in the 94\\AA\\ AIA band, so concluded that the hot emission seen in the  AIA 94~\\AA\\ band is effectively due to very hot plasma (6-8MK). More direct evidence has been found from observations of another Fe XVIII line by the  SUMER spectrometer on board the SOHO mission \\citep{2012ApJ...754L..40T}.   However, even if the indication is rather strong, it is still not enough to establish that the plasma is actually so hot, since in theory it is possible that plasma at a lower temperature, but with very high emission measure, can give the same line intensity, as suggested by \\cite{2012ApJ...754L..40T}. Indeed  \\cite{2013A&A...558A..73D} used simultaneous EIS and AIA observations of active regions cores to show that a significant fraction of the Fe XVIII 94~\\AA\\ intensity can be due to plasma  at 3~MK and not 6~MK. In fact,  \\cite{2013A&A...558A..73D} showed that often Fe XVIII 94~\\AA\\ emission is present in the cores of active regions, but Ca XVII (which has a narrower formation temperature, hence is sensitive to hotter plasma) is not. The only way to disentangle the various contributions to the AIA 94~\\AA\\ band  is therefore  to perform an emission measure modelling. To this purpose, here we use the same observations from Hinode and SDO as in Paper I, that include both high-resolution spectroscopic data, over a wide spectral window, and images with high spatial resolution.  All this information provides simultaneous constraints on the plasma thermal structure along the line of sight. To further support the analysis we replicate the same analysis  on the same data set but taking a region outside of the core that shows no evidence of these very hot components.We also compare two different inversion methods. In Section~\\ref{sec:obs} we describe the observation and the data analysis, and in Section~\\ref{sec:discuss} the results are discussed. ",
        "conclusions": "\\label{sec:discuss}  We analysed the thermal distribution of the plasma along the line of sight in two different regions where the narrow-band imagers diagnosed very different dominant temperatures, a hotter and a cooler one. The analysis was mostly based on the spectral data from Hinode/EIS. Our major attention was to the hottest components that may be signature of impulsive heating at work, and here our aim was to reconstruct the whole EM distribution and check the coherence of the overall scenario, including the hot component. At variance from many other similar analyses, our approach was to consider a limited number of spectral lines ($\\sim 1$ per temperature bin), and we picked the most intense lines uniformly covering the temperature range. In this way we simplified our control and interpretation of the reconstruction results, and all measured fluxes had the same weight in determining the global emission measure distribution. We also had a particular approach regarding the assessment of the uncertainties. The typical choice is to associate the same percentage error to all measured fluxes; 20\\% is the most commonly assumed value (e.g. \\citealt{2011ApJ...734...90W}). This uncertainty safely includes possible systematic unknown effects due to the instrument, atomic data and chemical abundances, but weights the same all measured values, independently of whether a line is strong or weak. Our choice is different. The measured fluxes were assigned exclusively the poissonian error, dictated only by the photon statistics. This allowed us to weight more the intense lines and to minimise the uncertainties of the solutions of the DEM reconstructions.  As a consequence, the DEM reconstruction with the MCMC technique leads to solution clouds with a narrow distribution around the best solution for many temperature bins (factor 2-3), significantly narrower than in previous works. Only a minority of temperature bins show a spread solution cloud. The constraint on the hot components is tighter.  The error estimate is important in this analysis, and we are aware that the real uncertainties are surely larger than those that we assume. The uncertainties can influence considerably the error on the DEM solution, and sometimes even the solution itself, as thoroughly discussed in \\cite{2012ApJ...745..111T}. We should also consider that the cloud spread may underestimate the real error on the DEM solution. The small uncertainties typically lead to a larger number of iterations with the MCMC method for better convergence \\citep{2012ApJ...745..111T}, but our 400 iterations is certainly an appropriate number for our cases. Tests show that DEM structures are reliably recovered on scales larger than $\\Delta \\log T = 0.2$, so we do not discuss narrower features.  The good agreement between predicted and observed line intensities confirms this, although we note that some significant discrepancies are present, in particular with the  Ca XV line, as also found previously \\citep{2013A&A...558A..73D}. The good agreement between the two inversion methods, already found in \\cite{2011A&A...535A..46D}, is confirmed. This suggests that the main source of uncertainty resides in the choice of parameters, atomic data and instrument calibration.  Although our assumptions probably underestimate the errors, the coherent support from other instruments (AIA, XRT), the agreement with the results from another reconstruction method and the coherence with the morphology seen in the images spanning the different temperature regimes makes us confident that other errors should not affect our results considerably.  Overall, the analysis has confirmed the two general characteristics anticipated by the imagers. The comparison of the DEM distributions of the cold and hot regions has, on the one hand, revealed substantial thermal components for $logT < 6.3$ in the cold region, without showing prominent features. The reconstruction of the cold region also shows  minor components for $logT > 6.3$. Since the images show no bright features in the hot channels and lines, we may use the values of these emission measure components as sensitivity limits of our analysis, i.e. we may not trust emission measure components below $10^{26}$ cm$^{-5}$.  The hot region shows a much more peaked thermal structure, with the positive gradient typically found in previous studies in coronal loops of active region cores \\citep{2011ApJ...734...90W}. The peak is at a rather high temperature ($\\log T = 6.6$) and beyond that the emission measure declines, but not very steeply and still showing a significant fraction of emission measure at temperature $\\log T \\leq 6.8$. Some components might be present at even higher temperature, although with higher uncertainties, up to the limit of the thermal sensitivity of our analysis ($\\log T \\sim 7$). We believe that the joint use of hot spectral lines, AIA 94 \\AA\\ channel and XRT filterbands, helps to partially remove the so-called \"blind spot\" for $\\log T > 6.8$ \\citep{2012ApJ...746L..17W}. Although the presence of the very hot component looks confirmed, still it is based on a very limited amount of information, and therefore some care should still be used. Some further feedback is provided by the comparison with the images, and between the images. The morphology of the region in the Ca XVII line only partially overlaps the morphology in immediately cooler lines (Fe XVI), while it matches well the X-ray and AIA 94 \\AA\\ morphology. This confirms that in the hot region selected in the present analysis there is indeed significant emission at 6~MK, resulting in strong  Ca XVII, so that the Fe XVIII emission in the  AIA 94 \\AA\\ band is also due to this component, unlike other cases in the cores of active regions where Ca XVII is not observed, and the  Fe XVIII emission is due to a large emission measure at 3~MK \\citep{2013A&A...558A..73D}.  Our analysis has remarked how spectral data are by far more constraining than the data from imagers, because the spectral lines are much more sensitive to temperature variations that the broader bands of the imagers (in agreement with the findings of \\cite{2012ApJ...758...54T}. Still it stresses that a quantum leap in the diagnostics of the hottest DEM components needs constraints from more lines sensitive to emission from high temperature plasma. These might be, for instance, easily accessible to broad-band X-ray spectrometers, to which we look forward in future space missions."
    },
    "1402/1402.0691_arXiv.txt": {
        "abstract": "Ultraviolet observations are the best tool to study hot stellar populations which emit most of their light at short wavelengths. As part of a large project devoted to the characterization of the UV properties of Galactic globular clusters (GGCs), we collected mid/far UV and optical images with the WFPC2@\\hst\\ for 31 GGCs. These clusters cover a wide range in metallicity and structural parameters, thus representing an ideal sample for comparison with theoretical models. Here we present the first results from an ongoing analysis aimed at deriving the temperature distribution of Horizontal Branch stars in GGCs. ",
        "introduction": "\\label{sec:intro} In old stellar populations like Galactic globular clusters (GGCs), the hottest stars are Blue Stragglers (BSS), Horizontal Branch (HB), Post-Early Asymptotic Giant Branch (PEAGB) and AGB-manqu\\'e whose spectral emission peaks at effective temperatures higher than $\\sim 6,000\\,K$ \\citep{Fer12,Moe07,Fer97,Swe87}. Therefore these stars mostly emit in the ultraviolet (UV) part of the electromagnetic spectrum \\citep{Moe95,deB85}. Since the Earth's atmosphere is almost completely opaque to this wavelength regime, any campaign for the observation of hot stars in UV must be performed using space facilities like, for instance, the Hubble Space Telescope (\\hst) or the Galaxy Evolution Explorer (GALEX). Unfortunately, the theoretical models describing the evolution of stars of the earliest spectral types still suffer from many uncertainties which are also due to the absence of a complete and homogeneous sample of UV emitters. Therefore, the UV observation of hot stars in GGCs is extremely important in order to compare data and evolutionary models in observational planes not hampered by critical color - effective temperature (\\teff) transformations \\citep{Dal13,Sch12,Dal11,Fer06-2,Cas04,Fer03,Fer98,Fer97-2}.  One of the most interesting stellar evolutionary sequences that still wait to be fully understood is the HB. It is commonly accepted that metallicity is the first parameter driving the Horizontal Branch morphology. Indeed, metal-rich GGCs typically have red HBs while metal-poor ones have more extended and blue HBs. However, there are several clusters with the same metallicity showing remarkable differences in Horizontal Branch morphology: nice examples are the couples \\ngc6388 - \\ngc5927, M\\,3 - M\\,13 and M\\,15 - M\\,92. Therefore, metallicity alone is not able to explain the observation of the complex HB zoology in GGCs \\citep{Fre81}. This issue, known as the `2nd parameter' problem, has focused the attention of several authors in the last years \\citep{Gra10,Dot10,Cat09,Lee94,Fus93}, but the solution is not obvious because many mechanisms play a role in shaping the color distribution of HB stars, such as mass loss, age and He abundance \\citep{Roo73}. Nonetheless, now there is a general consensus about the fact that age is the main global 2nd parameter \\citep{Gra10,Dot10}.  \\par The HB temperature extension can be an efficient way to parametrize the HB morphology and therefore to study the `2nd parameter' problem. This approach has been already attempted by \\citet{Rec06} and, later on, by \\cite{Bab09}, who used optical CMDs for their analysis. However, temperature derived from optical filters only, can be underestimated by more than $10,000\\,K$ in the case of very extended HBs \\citep[see the case of \\ngc2808, ][]{Dal11}. This demonstrates the difficulties of deriving properties of hot HB populations from pure optical data. Indeed in the optical plane, stars move down to the blue HB tail either because the bolometric correction increases or because stellar luminosity decreases. For very blue HB stars (like Blue Hook stars) both the above quantities are changing and optical colors cannot provide a reliable estimate of \\teff.  \\par As a part of a large project aimed at characterizing the UV properties of GGCs \\citep{San14,Dal13-2,Dal13,Con12,Dal12,San12,Sch12,Dal11,Dal09} we present here some preliminary results coming from the photometric analysis of UV observations of 31 GGCs obtained with the Wide Field and Planetary Camera 2 (WFPC2) and finalized to accurately derive the HB temperature distribution and extension. ",
        "conclusions": "We have analyzed far/mid UV-optical images for ten out of 31 GGCs collected within the WFPC2 UV survey. The main task of this analysis is to derive the extension and the temperature distribution of the HB of the targeted clusters taking advantage of the UV data. The use of UV bands is the best tool to accomplish our task since, for stars with high \\teff, the bolometric correction decreases and the stellar luminosity increases with respect to the optical bands. It has been shown, indeed, that the temperatures of extended HBs derived from a combination of pure optical filters can be underestimated by $10,000 - 15,000\\,K$ \\citep{Dal11}. We used \\ngc6293 as a starting point in our analysis, determining the temperature distribution of its HB stars. We are now applying this analysis to other GGCs in our sample with extended HBs, like $\\omega\\,Cen$ and M\\,15."
    },
    "1402/1402.0002_arXiv.txt": {
        "abstract": "Disc-winds originating from the inner parts of accretion discs are considered as the basic component of magnetically collimated outflows. The only available analytical MHD solutions to describe disc-driven jets are those characterized by the symmetry of radial self-similarity. However, radially self-similar MHD jet models, in general, have three geometrical shortcomings, (i) a singularity at the jet axis, (ii) the necessary assumption of axisymmetry, and (iii) 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, by solving the full three-dimensional equations of MHD and impose a termination radius at finite radial distance. We focus here on studying the effects of relaxing the (ii) assumption of axisymmetry, i.e. of performing full 3D numerical simulations of a disc-wind crossing all magnetohydrodynamic critical surfaces. We compare the results of these runs with previous axisymmetric 2.5D simulations. The structure of the flow in all simulations shows strong similarities. The 3D runs reach a steady state and stay close to axisymmetry for most of the physical quantities, except for the poloidal magnetic field and the toroidal velocity which slightly deviate from axisymmetry. The latter quantities show signs of instabilities, which, however, are confined to the region inside the fast magnetosonic separatrix surface. The forces present in the flow, both of collimating and accelerating nature, are in good agreement in both the 2.5D and the 3D runs. We conclude that the analytical solution behaves well also after relaxing the basic assumption of axisymmetry. ",
        "introduction": "Astrophysical jets are observed in association with a wide range of objects, from Brown Dwarfs and young stellar objects to supermassive Black Holes in active galactic nuclei; however, there are still open several questions concerning the launching and acceleration mechanisms of jets. In all cases, jets and discs are inter-related, while magnetic fields play a key role in accelerating the outflows. \\citet{BlP82} studied the magneto-centrifugal acceleration along magnetic field lines threading an accretion disc. They showed the braking of matter in the azimuthal direction inside the disc and the outflow acceleration above the disc surface guided by the poloidal magnetic field components. Toroidal components of the magnetic field then collimate the outflow. Numerous semi-analytic models extended the work of Blandford \\& Payne along the guidelines of radially self-similar solutions of the full magnetohydrodynamic (MHD) equations \\cite[e.g.][]{CoL94, Li95, Li96, Fer97, VlT98}. Several numerical studies exist, which have focused on the magnetic launching of disc-winds. In most models a polytropic equilibrium accretion disc has been used as a boundary condition \\citep[e.g.][]{UKR95, KLB99, KLB03, OCP03, NaM04, ALK05, ALK06, PRO06}. The magnetic feedback on the disc structure was therefore not calculated self-consistently. Only in recent years have been carried out the first simulations which include self-consistently the accretion disc in the calculations of jet launching \\citep[e.g.][]{CaK02, CaK04, KMS04, ZFR07, TFM09, MFZ10, SFP12, TFM13, FeS13}.  Often analytical solutions and their integrals of motion are used for connecting observed quantities far away from the jet source with properties of the jet-driving accretion disc. The behaviour of integrals of motions has been also tested in numerical simulations. Several 2.5D simulations of disc winds from the disc as a boundary condition \\citep[e.g.][]{RUK97, OuP97, KLB99, UKR99} found super-fast magnetosonic flows with properties which are expected from self-similar theory, although they have not used self-consistent analytical solutions for describing the boundary conditions. Recently, \\citet{SNO10} performed 3D simulations extending the calculations of \\citet{OuP97} and found good agreement between both. In addition, \\citet{SNO10} also simulated configurations closer to self-similar solutions (their run BP) and compared synthetic emission maps derived from them with HST observations. A similar comparison is presented in \\citet[][hereafter paper~II]{SGT10} using simulations from \\citet[][hereafter paper~I]{STV08}.  Most self-similar models have three serious limitations, (1) the flow often does not cross all critical points (especially not the fast-magnetosonic limiting characteristic), with the result that the terminal wind solution is not causally disconnected from the disc, (2) singularities exist at the jet axis in radial self-similar models, and (3) for deriving self-similar models the assumption of axisymmetry is necessary.  \\citet[][V00 hereafter]{VTS00} showed that a terminal wind solution can be constructed that is causally disconnected from the disc and hence any perturbation downstream of the superfast transition cannot affect the upstream structure of the steady outflow. The other two limitations can only be solved using numerical simulations extending the analytical solution of e.g. V00 as done by \\citet[]{GVT06} using the MHD code NIRVANA \\citep{Zie98}, \\citet[]{MTV08} using the MHD code PLUTO\\footnote{publicly available at http://plutocode.ph.astro.it/} \\citep{MBM07} and again with PLUTO in paper I for comparison with models with finite jet-emitting disc radii. \\citet{CGV08} extended the solution by adding the effects of resistivity. However, all of those models still assumed axisymmetry.  The aim of this paper is to investigate numerically how relaxing the assumption of axisymmetry, i.e. performing full 3D numerical simulations of a disc-wind crossing all magnetohydrodynamic critical surfaces, affects the topology, structure and stability of this particular radial self-similar analytical solution.  This paper is organized as follows: the numerical setup is briefly described in Sec.~\\ref{sec_num_models}. In Sec.~\\ref{sec_results} we describe the results of our simulation. We close with a summary and discussion of the results in the last section. ",
        "conclusions": "We have shown the results of the first 3D simulation of a disc-wind crossing all magnetohydrodynamic critical surfaces. This result is important in order to assure that the main flow is causally disconnected from its source, which is a prerequisite for a jet with the observed long lifetime. We have compared these results with previous axisymmetric 2.5D simulations.  The structure of the flow in all simulations exhibits strong similarities. In the outer part of the flow, its structure is almost identical in all cases. Near and at the position of the fast magnetosonic separatrix surface (FMSS) which shows as a shock, some minor deviations of the 3D run from the 2.5D runs are present.  The 3D runs reach a steady state and stay close to axisymmetry for most of the variables, except for the poloidal magnetic field and the toroidal velocity which deviate considerably from axisymmetry, but are not dynamically important. The latter quantities show signs of instabilities, which, however, are confined to the region inside the FMSS. It is important to emphasize that the crossing of the FMSS, which represents the ``event horizon'' for the propagation of MHD waves, does not allow any disturbance at large distances to reach the base of the flow, contributing to its stability.  The forces present in the flow, both of collimating and accelerating nature, are in good agreement in both the 2.5D and the 3D runs.  The main goal of the present paper is to check whether the axisymmetric, radially self-similar MHD solution for a polytropic disc-wind which crosses all appropriate critical surfaces to satisfy causality \\citep{VTS00}, behaves ``well'' also after (i) removing the axial singularity and (ii) relaxing the assumption of axisymmetry.  We have found that this solution is structurally stable to non-axisymmetic perturbations. The next step will be to further configure and improve this solution such as it may describe realistic astrophysical disc-winds. As in paper~I, this further extension will involve the truncation of the solution at some arbitrary radii and use of this truncated solution to study the temporal evolution of jets within an MHD model and further comparison with observations. \\citet{SNO10} argue that models exhibiting a slowly varying poloidal field component in the accretion disc (their model OP) match observations better than those resembling self-similar MHD solutions (as their model BP). Note, however, that in paper~II we showed that radially truncated self-similar solutions may very well describe real HST observations."
    },
    "1402/1402.0833_arXiv.txt": {
        "abstract": "In the theory of resonant mode coupling, the parent and child modes are directly related in frequency and phase.  The oscillations present in the fast rotating $\\delta$\\,Sct star KIC\\,8054146 allow us to test the most general and generic aspects of such a theory. The only direct way to separate the parent and coupled (child) modes is to examine the correlations in amplitude variability between the different frequencies.  For the dominant family of related frequencies, only a single mode and a triplet are the origins of nine dominant frequency peaks ranging from 2.93 to 66.30 cycles day$^{-1}$ (as well as dozens of small-amplitude combination modes and a predicted and detected third high-frequency triplet).  The mode-coupling model correctly predicts the large amplitude variations of the coupled modes as a product of the amplitudes of the parent modes, while the phase changes are also correctly modeled. This differs from the behavior of ``normal'' combination frequencies in that the amplitudes are \\emph{three orders of magnitude} larger and may exceed even the amplitudes of the parent modes.  We show that two dominant low frequencies at 5.86 and 2.93 cycles day$^{-1}$ in the gravity-mode region are not harmonics of each other, and their properties follow those of the almost equidistant high-frequency triplet. We note that the previously puzzling situation of finding two strong peaks in the low-frequency region related by nearly a factor of two in frequency has been seen in other $\\delta$ Sct stars as well. ",
        "introduction": "\\begin{figure*}[!t] \\centering \\includegraphics[bb=40 70 750 550, width=\\textwidth, clip]{Fig1.eps} \\caption{The dominant 9 frequencies of the T family of frequencies. The frequency numbers, e.g. f26, denote their approximate values in cycles day$^{-1}$ and this notation will be used in the rest of the paper. Note that the frequencies are both in the low-frequency g-mode and high-frequency p-mode regions. The bottom panel illustrates their relation to all the detected frequencies.} \\end{figure*}  The recent spaced-based studies of $\\delta$\\,Sct and $\\gamma$\\,Dor stars have revealed that most of these stars pulsate with a multitude of pulsation modes in both the low and high-frequency regions (Uytterhoeven et al. 2011), where the gravity (g) and pressure (p) modes are found. The numerical relationships between many low and high frequencies indicate that it is not possible to treat these regions independently of each other. For example, the low-frequency region contains g-mode pulsations as well as rotation peaks (Balona 2011, Breger et al. 2011), while the high-frequency region contains p-modes as well as high-degree g-modes (Monnier et al. 2010, Breger et al. 2013). Regrettably, pulsation mode identifications are generally not available for these fainter stars studied from space; consequently, frequency patterns have been the most valuable tool. In the present investigation, we add amplitude and phase variations as an additional method to determine the modal origin.  The rapidly rotating ($V \\sin i = 300\\,$km/s) hybrid $\\delta$\\,Sct/$\\gamma$\\,Dor pulsator KIC\\,8054146 has been measured extensively by the {\\it Kepler} spacecraft. Altogether, 349 statistically significant frequencies were determined from three years of short-cadence {\\it Kepler} data (Breger et al. 2012, hereafter Paper 1). The amplitudes ranged from about 2 ppm (parts-per-million) to 200 ppm.  The excellent frequency resolution of three years of data for KIC\\,8054146 also revealed that many frequencies are related over and beyond the expected simple harmonics and combinations. In fact, three separate families of frequencies spanning a 200 cycles day$^{-1}$ frequency range have been discovered so far. Within each family, the amplitude variations of the low-frequency members correlate with those of the high-frequency members. The unprecedented accuracy of the {\\it Kepler} allows us examine the physical origin of these families in detail. One of these families could already be interpreted as high-degree prograde Kelvin modes (Breger et al. 2013).  \\begin{table} \\caption{Dominant frequencies of the ``T'' family} \\begin{tabular}{llrl} \\hline \\noalign{\\smallskip} Name & \\multicolumn{2}{c}{Frequency} & Comment\\\\ & cycles day$^{-1}$ & $\\mu$Hz \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} f3 & 2.9308 & 33.92 & Low frequency\\\\ f3b & 2.9334 & 33.95 & Low frequency\\\\ f6 & 5.8642 & 67.87 & =f3+f3b\\\\ f26 & 25.9509 & 300.36 \\\\ f34 & 34.4836 & 399.12 & Triplet 2\\\\ f37 & 37.4170 & 433.07 & Triplet 2\\\\ f40 & 40.3479 & 466.99 & Triplet 2\\\\ f60 & 60.4346 & 699.47 & Triplet 1\\\\ f63 & 63.3680 & 733.43 & Triplet 1\\\\ f66 & 66.2988 & 767.35 & Triplet 1\\\\ \\hline \\end{tabular} \\end{table}  This paper investigates the dominant set of modes excited in KIC\\,8054146, which we have called the T family. The tools for this investigation are not only the frequency values, but also their large and steady amplitude and phase changes over the three years.  The bottom panel of Fig. 1 shows the amplitude spectrum of KIC 8054146 in the 0 to 100 cycles day$^{-1}$ range. For more observational information we refer to Paper I, which was extended to include the {\\it Kepler} short-cadence data from quarters 11 to 14 (Q11 to Q14). Since the frequency resolution of the combined Q5-Q14 data exceeds the time scales of the (small) frequency variations of the star, the diagram actually shows the sum of four subsets. The top panel of Fig. 1 presents the T family of dominant modes with additional properties of these modes listed in Table 1.  While equidistant frequency spacings and a numerical relationship between the low and high frequencies are also seen in other $\\delta$ Sct stars studied by $\\it Kepler$, the number and sizes of such occurrences in KIC 8054146 is unusual. In addition, the continuous short-cadence coverage over three years makes a study of the correlations of amplitude variability of the different modes possible: this allows us to uniquely separate the parent and child (coupled) modes. ",
        "conclusions": "For KIC 8054146 we identify coupled modes and can uniquely distinguish which are the parent and child modes. This is made possible by the fact that for the simplest, lowest-order mode coupling the amplitudes of the coupled modes are given by the product of the amplitudes of the parents. We speculate that the simple mode-coupling model is applicable due to the relatively small amplitudes and therefore lowest-order nonlinear interaction of the modes involved. Consequently, the present study is only made possible by the excellent time coverage and precision of the data from the {\\it Kepler} spacecraft. We note that the generally unknown mode-coupling constants do not present a problem for KIC 8054146 due to the large amount of amplitude variability over the three years of data; a single short stretch of data allows us to derive the mode-coupling coefficients, which subsequent data groups show to be constant.  The state of nonlinear or combination frequencies detected in KIC\\,8054146 is the following: \\begin{enumerate} \\item We detect ``normal'' nonlinearities in this star; these are combination frequencies for which the amplitudes are related by $A_3\\sim  \\mu * A_1 * A_2$, where $\\mu$ is $\\sim 4$. These small nonlinearities are commonly seen in most if not all Delta Scuti stars.  \\item KIC 8054146 also possesses larger-amplitude nonlinear signals, and we are able to use the measured amplitude and phase variations of these signals to determine which modes are the ``parent'' modes and which are the ``child'' modes. The derived value of $\\mu$ relating the  amplitudes of child and parent modes is $\\mu \\sim1000$ to 10,000 or more. \\end{enumerate}  A generic nonlinear model (such as the nonlinear relation between temperature and flux at the surface) is not able to simultaneously produce both these ``small'' and ``large'' nonlinearities. On the other hand, if a subset of the combination frequencies lie close to eigenfrequencies in the star, then these modes and only these modes can be resonantly driven to larger amplitudes. While this is far from the last word on the subject, the data are fully consistent with this ``resonant mode coupling'' scenario.  In particular, the present investigation showed that in KIC 8054146, two dominant low frequencies are actually coupled modes with amplitudes larger or equal to those of the dominant parents in the higher-frequency domain. The low-frequency region of these so-called hybrid pulsators contains a large number of linearly driven gravity modes, as was confirmed by Chapellier et al. (2012) from searching for equidistant period spacings in ID 105733033. Furthermore, in a number of hybrid pulsators, some frequency values in the gravity-mode and pressure-mode domains are numerically not independent of each other. Given the existence of these gravity modes, we speculate that these could be the ``child'' modes that are resonantly excited to larger amplitudes. A confirmation of this resonance explanation would be provided if additional data could reveal that these particular low frequencies are part of a sequence of low-frequency modes equidistantly spaced in period."
    },
    "1402/1402.3479_arXiv.txt": {
        "abstract": "We present a study of age-related spectral signatures observed in 25 young low-mass objects that we have previously determined as possible kinematic members of five young moving groups: the Local Association (Pleiades moving group, age=20 - 150 Myr), the  Ursa Major group (Sirius supercluster, age=300 Myr), the Hyades supercluster (age=600 Myr), IC 2391 supercluster (age=35--55 Myr) and the Castor moving group (age=200 Myr). In this paper we characterize the spectral properties of observed high or low resolution spectra of our kinematic members by fitting theoretical spectral distributions. We study signatures of youth, such as lithium~{\\sc i} 6708 \\AA, H$\\alpha$ emission and other age-sensitive spectroscopic signatures in order to confirm the kinematic memberships through age constraints. We find that 21 (84\\%) targets show spectroscopic signatures of youth in agreement with the age ranges of the moving group to which membership is implied. For two further objects, age-related constraints remain difficult to determine from our analysis. In addition, we confirm two moving group kinematic candidates as brown dwarfs. ",
        "introduction": "Identifying members of known moving groups (MG) or open clusters can provide an important constraint on their age and composition. Low-mass stars and brown dwarfs (BDs) members of known MGs or open clusters can provide important feedback to atmospheric and evolutionary models of such objects at young ages, which currently await good calibration from constraints on age and composition. Therefore, there exists an imperative to compile a sample of MG members in this mass range that can be used as anchor points or test-beds for these models.  In this paper we present the last in a series of results of a major survey of Ultra-Cool Dwarfs (UCDs; objects with a spectral classification of M7 or later corresponding to Teff$<$2500 K), in young moving groups which began with the results presented in Clarke et al. (2010), hereafter Paper I, and G\\'alvez-Ortiz et al. (2010), hereafter Paper II. We focus here on the study of age-sensitive spectroscopic signatures such as gravity sensitive features (e.g. Gorlova et al. 2003; McGovern et al. 2004), H$\\alpha$ emission (e.g. West et al. 2004, 2008, 2011) and the presence of Li~{\\sc i} 6708 \\AA \\ in a sample of objects previously classified as members of a MG based on kinematic or astrometric criteria.  Despite the recent disagreement about the origin of moving groups (Famaey et al. 2007, 2008; Antoja et al. 2008; Klement et al. 2008; Francis \\& Anderson 2009; Zhao et al. 2009; L\\'opez-Santiago et al. 2009; Bovy \\& Hogg 2010; Murgas et al. 2013), and as we discussed in Paper II, we are assuming the classical concept: a moving group is a young stellar population that shares a common space motion (e.g. Pinfield et al. 2006) whose members have a common origin, and therefore, age and composition. We focused our studies on well-documented groups: the Hyades supercluster (HY; 600 Myr), the Ursa Major group (Sirius supercluster) (Si; 300 Myr), IC 2391 supercluster (IC; 35-55 Myr), the Castor Moving Group (CA; 200 Myr), and the Local Association (20-150 Myr) or Pleiades (PL) moving group (See Paper I and II for details of MG properties and references thereof).  Many efforts have been made to obtain accurate models to understand the cool and complex atmospheres of low mass stellar and substellar objects (e.g. Kirkpatrick et al. 1993; Rajpurohit et al. 2011; Reyle et al. 2011; Rajpurohit et al. 2012). But the models still show discrepancies, for example in the strength of some absorption bands: discrepancies likely due to inaccurate atomic parameters and/or missing molecular opacities. Also, dust formation and its behaviour with temperature changes spectral characteristics in many ways. Models that describe Very Low-Mass stars (VLMS) and BD atmospheres cannot hope to be in agreement with observations without the inclusion of dust. The onset of iron and silicates dust grain formation suspended in the photospheres of late-type M dwarfs through the L dwarf spectral sub-types is accompanied by spectral reddening, especially in near-IR (based upon the size, quantity and distribution of this dust). The T dwarf spectral sub-type sequence demonstrates a reversal to bluer colours due to the dust settling below the photosphere, thus reducing the reddening. Consequently, atmospheric models that describe dwarf objects with T$_{\\rm eff}$ $\\leqslant$ 2800 K should include treatment for the effects of atmospheric dust evolution.  Up to date, there are several families of stellar model atmospheres of very low mass stars. These models are still incomplete or approximate in some physical properties (opacities, oscillator strengths for some lines and molecular bands, etc), showing few or considerable discrepancies depending on the atmospheric regions, temperatures, etc, and there are also non covered areas or \"gaps\". A comprehensive study of models tested by observations is needed to fully understand this cool atmospheres. Since BDs are objects occupying an intermediate position between stars and giant planets, studying their atmospheres can help us to better understand the processes in giant exo-planet atmospheres. The determination of the physical properties of VLMS and BDs is also important for understanding a broad range of topics including stellar and planetary formation, circumstellar disks, dust formation in cool atmospheres, and the initial mass function.  By comparing high and low resolution optical spectra of our 25 low-mass stars and BD kinematic MG candidates with a grid of atmospheric models, we can obtain effective temperature and surface gravity estimates.  In Section~2 we describe the sample characteristics and selection criteria, while in Section~3 we present details of our observations and data reduction techniques. In Section~4 we explain the model atmosphere fitting process, and in Section~5 we describe the age indicators and the analysis used to assess the sample memberships. Finally, in Section~6 we present a brief discussion and summary of our results. ",
        "conclusions": "We have presented a study of the spectral signatures of 25 low-mass objects that were candidate members of five young moving groups. We studied different typical age-constraining spectroscopic criteria utilising high and low resolution spectra, and combined the results to extract a final membership assessment for each target.  We took into account spectral classification and thus approximated atmospheric temperature to apply the most appropriate model or models in the search of the best physical parameter constraints. When we used the synthetic stellar atmospheric models, the good agreement between observed and theoretical energy distributions suggests that semi-empirical models describe well the impact of dust in the atmospheres of M dwarfs. $S$-function analysis shows that the inclusion of dust effects makes it possible to achieve better fits for objects with $T_{\\rm eff}\\leqslant2800$~K.  In both samples, although youth can generally be established, $log~g$ and $EW$(Na~{\\sc i}) were not useful for discriminating between different MGs. Some targets of sample A that were candidates to the HY moving group having lower $log~g$ values than candidates to younger MGs, and similarly in $EW$(Na~{\\sc i}) values. In sample B the values are more consistent but most of the targets have several candidatures, making the situation more complicated. However, looking at the final membership results, only one target of sample A is classified as a member of a MG with age $<$200 Myr, which makes this criterion non applicable for most of them. For sample B, in spite of the multiple membership candidatures in the final results, we can discriminate between the lower values of $log~g$ (or lower values of $EW$s(Na~{\\sc i})) for targets that are members of younger MGs, medium values for targets that are members of intermediate age MGs and the higher values for targets that are members of oldest MGs, which finally is consistent with the $log~g$-age relation. The comparison of $EW$(Na~{\\sc i})s is probably useful to discriminate very young targets (less than 100 Myr) in a homogeneous sample (as suggested by Schlieder et al. 2012) while $log~g$ values are probably useful to discriminate very young (less than 100 Myr) and very old (more that 600 Myr) objects, serving as complementary study to other criteria.  The activity-age relation is a reasonably useful age criterion up to $\\approx$M7 spectral type but activity variability, binarity, and the comparison of different instruments/resolutions can introduce considerable uncertainties in age estimation. The presence of Lithium discriminates between young and old low mass stars but an ambiguity between young low mass stars and BDs limits its use in this cool temperature region. Gravity sensitive features and rotational velocity can be useful youth indicators in support of activity and lithium diagnostics, but clearly have large dispersion and they need to be applied in conjunction with similar data. For a reliable constraint of cool object ages, a combination of different criteria are needed.  We find 10 of the 13 objects from sample A to be probable members of one of the MGs, and also that 2MASS0020-2346 is a young disk member. SIPS0007-2458, a candidate member of the IC2391 MG, shows positive candidature under most criteria but does not show the expected lithium (associated with the age of the IC2391 MG). Due to the age constraint from $H\\alpha$ emission and surface gravity, we are inclined to think that SIPS0007-2458 is probably a young object from the YD class, but is not in the IC2391 MG, although it could also be a contaminant field object. 2MASS0334-2130 is a similar case, with an M4.5-6 spectral type, it should contain lithium if it were an IC2391 MG member. Except for the $H\\alpha$ criterion, that could be reflecting a maximum in activity level or the influence of an unseen companion, everything indicates that 2MASS0334-2130 is older than the IC2391 MG. Thus we also conclude that 2MASS0334-2130 is not an IC2391 member, although it could be a young object from the YD, or a contaminant target from the field.  In sample B, 10 objects could be MG members, although more information is needed to discriminate between the possible MG membership of objects with several candidatures. We could not obtain any results for 2MASS1909-1937, so this remains an astrometric candidate to the Castor MG. 2MASS1734-1151 gave contradictory results so it also remains an astrometric Pleiades MG candidate until further analysis can be performed.  We also confirm two moving group candidates, SIPS2045-6332 and SIPS2039-1126, as BDs.  With all the acquired information, we find that 85\\%, 83\\% and 84\\% of samples A, B, and both combined, show spectroscopic signatures of youth in agreement with the age of the moving group to which they present kinematic membership. This result suggests that there is a fairly low rate of contamination in such kinematic candidate samples. Also, there are candidates that cannot be confirmed or dismissed with the information obtained, which might (in the future) increase the final confirmation rate. We note that additional diagnostics to assess membership may be measured in the future, such as chemical tagging. Recently, in a kinematical-chemical investigation of the AB Dor moving group \u201cStream\u201d, Barenfeld et al. (2013) shows that kinematics, color-magnitude positions, and stellar youth indicators alone can be insufficient for testing whether a kinematic group of stars actually shares a common origin. Future chemical study of our MG candidates will complement this study.  In addition, this study has also helped to test and improve the atmosphere models for cool dwarfs (see Kuznetsov et al. 2012 and Kuznetsov et al. 2013). Youth can be an advantage for various reasons. Because young objects are brighter they can be detected and studied out to greater distance than older, fainter objects. Also, the proximity of the MGs can allows the exploration of the faint circumstellar environment at relatively small distances from the star. Indeed, we will target our strongest MG candidates to search for lower mass companions using high resolution imaging techniques (e.g. adaptive optics). The discovery of planetary systems around young and low-mass objects provides crucial information for the understanding of planetary and stellar formation.  From our complete study of the 80 objects considered (within papers I, II and this paper), we have found a total of 45 new possible MG members (with 8 of them showing more than one MG candidature) and around 21 possible young disk objects with no clear membership of the 5 MGs considered.  Of these we find 26, 13, 6, 4 and 5 possible members of the Hyades, Castor, Ursa Major, Pleiades and IC2391 MGs respectively. Tables~\\ref{tab:comptod} and \\ref{tab:compothers} of the appendix compile all the M-L dwarfs investigated in our study.  \\begin{figure*} \\centering \\includegraphics[width=8.5cm,angle=0,clip]{Figura1_1.ps} \\includegraphics[width=8.5cm,angle=0,clip]{Figura1_2.ps} \\caption{Left: Observed SEDs of objects from sample A, the 13 targets plus reference objects in the 7600 to 8400 \\AA \\ range. Right: Observed SEDs of objects from sample B, the 12 targets plus one reference object in the 6500 to 9000 \\AA \\ range. They are ordered by spectral type derived from spectroscopy.} \\label{fig:todashl} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=6.0cm,angle=270,clip]{Figura2_1.eps} \\includegraphics[width=6.0cm,angle=270,clip]{Figura2_2.eps} \\caption{Sample A: random example of fits of theoretical spectra to the observed SEDs.} \\label{fig:fight} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=6.0cm,angle=270,clip]{Figura3_1.eps} \\includegraphics[width=6.0cm,angle=270,clip]{Figura3_2.eps} \\caption{Sample B: random example of the best fits to the observed SEDs. } \\label{fig:figlt} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=6.0cm,angle=270,clip]{Figura4_1.eps} \\includegraphics[width=6.0cm,angle=270,clip]{Figura4_2.eps} \\caption{Sample A: random example of the best fits to the spectral regions across K~{\\sc i} lines. } \\label{fig:fighk} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=6.0cm,angle=270,clip]{Figura5_1.eps} \\includegraphics[width=6.0cm,angle=270,clip]{Figura5_2.eps} \\caption{Sample A: random example of the best fits to the spectral regions across Na~{\\sc i} lines. } \\label{fig:fighna} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=7.0cm,angle=270,clip]{Figura7.eps} \\caption{ This figure (updated Figure~9 from Paper I), illustrates how an object would be selected for the various methods of age constraining. Candidates that appear to the right of the lithium edge (blue continuous line) can be followed up with a lithium test programme. Objects that appear younger than 200 Myr (horizontal red dashed line), are eligible for follow-up using spectroscopic gravity sensitive features. Vertical and horizontal orange dot-dashed line represents the Schlieder et al. (2012) limits for the Na~{\\sc i} diagnostic applicability. We here took into account the 200 Myr limit. Candidates that fall to the left left of the spectral type=M7 limit (vertical black dotted line), would thus be suitable for age/activity relation follow-up, although candidates with a spectral type close to M7 may be subject to large uncertainties on their age. Some candidate cannot be tested by any of these methods but will be eligible for age testing using $v\\sin{i}$. Sample A targets are plot as blue circles and sample B as red circles. We plotted up to two candidatures for targets with multiple MG.} \\label{fig:modelo} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=8.0cm,clip]{Figura8n.ps} \\includegraphics[width=8.0cm,clip]{Figura8nb.ps} \\caption{ The right and left panels show two versions of the same information with different legends in order to highligh different information. Left: From Mart\\'in et al. (2010), 65 high-gravity field objects are plotted as six pointed asterisk symbol, 6 low-gravity objects as open triangles, 12 reference field stars as open circles and 7 Upper Sco candidates as solid hexagons. We overplot sample A as blue crosses, and sample B as red filled circles. The red star marks LP944-20. Errors are approximately the size of plot symbols. Right: Here we overplot to Mart\\'in et al. (2010) data, sample A (crosses and filled squares), and sample B (filled circles and triangles), plotted in different colours depending on their final MG membership candidature (Table~\\ref{tab:resultfinal}). When objects present more than one candidature, we plotted the membership to the youngest MG. To better discriminate young targets, we use triangles and squares for candidates finally classified as possible member of MGs with ages $\\geq$200 Myr while crosses and circles are candidates finally classified as possible member of MGs with ages $<$200 Myr (see Sect. 6). } \\label{fig:ewnacomp} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=8.5cm,clip]{Figura9_1.ps} \\includegraphics[width=8.5cm,clip]{Figura9_2.ps} \\caption{ $EW$(H$\\alpha$) versus spectral type. In the same way as in Fig.~7, both panels show two versions of the same information with different legends in order to highligh different information. Left: Sample A is plotted as blue crosses, Sample B as red circles and literature known-age objects in different symbols according to age: red asterisk for objects with age between 1-12 Myr, blue diamonds for 40-50 Myr, violet asterisk for 90-100 Myr, blue asterisk for 120 Myr, red diamonds for 150 Myr, black asterisk for 300 Myr, black diamonds for 400 Myr and green asterisk for objects with ages over 1200 Myr. Literature data has been obtained from Terndrup et al. (2000), Mohanty \\& Basri (2003) and Shkolnik et al. (2009) where ages were calculated by known-age MG or Cluster membership or by other age constraining methods. Right: Similar to left panel where to favor the age discrimination we plot only three age intervals. Sample A targets are plotted as filled squares and sample B as filled triangles, where targets classified with ages $\\geq$300 Myr are plotted in blue and targets with ages $<$300 Myr are plotted in red. } \\label{fig:ewhacomp} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=6.0cm,angle=270,clip]{Figura10_1.eps} \\includegraphics[width=6.0cm,angle=270,clip]{Figura10_2.eps} \\caption{Fit to the observed spectra across Li~{\\sc i} resonance doublet for SIPS2045-6332 (left) and SIPS2039-1126 (right). Best synthetic fits are overplot with different Li abundances. } \\label{fig:lithium} \\end{figure*}   \\begin{table*} \\caption[]{A: Properties for sample A \\label{tab:samplea}} \\begin{flushleft} \\begin{center} \\small \\begin{tabular}{lllccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Name & $\\alpha$ (2000) & $\\delta$ (2000) & J-mag & Moving group$^{1}$ & $V\\sin{i}$$^{2}$  \\\\ & (h m s) & ($^{\\rm o}$ ' '') &  & candidature   & (km s$^{-1}$) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} SIPS0007-2458  &  0  7  7.800 & -24 58  3.80 &  13.11$\\pm$0.02 & IC     & 18 \\\\ 2MASS0020-2346 &  0 20 23.155 & \u221223 46  5.38 &  12.35$\\pm$0.02 & YD    & 12 \\\\ DENIS0021-4244 &  0 21  5.896 & -42 44 43.33 &  13.52$\\pm$0.03 & IC, CA & 11 \\\\ SIPS0027-5401  &  0 27 23.240 & \u221254  1 46.20 &  12.36$\\pm$0.02 & HY     & 27 \\\\ SIPS0153-5122  &  1 53 11.430 & \u221251 22 24.99 &  13.45$\\pm$0.03 & IC, HY & 14 \\\\ SIPS0214-3237  &  2 14 45.440 & -32 37 58.20 &  14.01$\\pm$0.02 & HY     & 18 \\\\ SIPS0235-0711  &  2 35 49.470 &  \u22127 11 21.90 &  12.45$\\pm$0.03 & HY     & 22 \\\\ 2MASS0334-2130 &  3 34 10.657 & \u221221 30 34.35 &  11.91$\\pm$0.02 & IC     & $<$10$^{a}$ \\\\ SIPS2039-1126  & 20 39 13.081 & -11 26 52.30 &  13.79$\\pm$0.03 & PL     & $>$15 \\\\ SIPS2045-6332  & 20 45  2.278 & -63 32  5.30 &  12.62$\\pm$0.03 & PL, CA & $>$15 \\\\ LEHPM4908      & 22 36 42.656 & -69 34 59.30 &  12.68$\\pm$0.02 & HY     & - \\\\ 2MASS2254-3228 & 22 54 58.110 & \u221232 28 52.20 &  13.58$\\pm$0.03 & PL, CA & 19 \\\\ LEHPM6542      & 23 57 54.822 & \u221219 55 1.89  &  13.31$\\pm$0.02 & HY     & 16 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^{1}$ MG to which targets are kinematic candidate (Paper II). HY= Hyades MG; SI= Ursa Major group; CA= Castor MG; PL= Pleiades; IC = IC 2391 MG; YD= Other young disk object.\\\\ $^{2}$ From paper II.\\\\ $^{a}$ Measused with low S/N.\\\\ \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{B: Properties for sample B \\label{tab:sampleb}} \\begin{flushleft} \\begin{center} \\small \\begin{tabular}{lllcccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Name & $\\alpha$ (2000) & $\\delta$ (2000) & J-mag & Moving group   \\\\ & (h m s) & ($^{\\rm o}$ ' '') & &  candidature   \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2MASS0814-4020 & 08 14 35.46 & -40 20 49.26 & 14.356$\\pm$0.023 & HY \\\\ 2MASS1146-4754 & 11 46 51.04 & -47 54 38.17 & 14.897$\\pm$0.042 & SI \\\\ 2MASS1236-6536 & 12 36 32.38 & -65 36 35.6  & 15.277$\\pm$0.057 & PL, CA, IC, SI \\\\ 2MASS1326-5022 & 13 26 53.48 & -50 22 27.04 & 14.715$\\pm$0.037 & IC, CA \\\\ 2MASS1433-5148 & 14 33 41.95 & -51 48 03.70 & 14.206$\\pm$0.034 & CA, PL, SI, IC \\\\ 2MASS1557-4350 & 15 57 27.39 & -43 50 21.47 & 14.224$\\pm$0.028 & PL, CA, IC \\\\ 2MASS1618-3214 & 16 18 08.92 & -32 14 36.17 & 14.920$\\pm$0.040 & IC, CA, SI \\\\ 2MASS1734-1151 & 17 34 30.53 & -11 51 38.83 & 13.110$\\pm$0.028 & PL \\\\ 2MASS1736-0407 & 17 36 56.09 & - 4 07 25.84 & 15.516$\\pm$0.070 & SI, CA \\\\ 2MASS1745-1640 & 17 45 34.66 & -16 40 53.81 & 13.646$\\pm$0.026 & SI, CA, HY \\\\ 2MASS1756-4518 & 17 56 29.63 & -45 18 22.47 & 12.386$\\pm$0.019 & CA \\\\ 2MASS1909-1937 & 19 09 08.21 & -19 37 47.96 & 14.520$\\pm$0.026 & CA \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Details of observing runs \\label{tab:obs}} \\begin{flushleft} \\begin{center} \\small \\begin{tabular}{ccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Number & Date & Telescope & Instrument & Spect. range & Dispersion    \\\\ &  &           &              & (\\AA)   & (\\AA) \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1 & 28/03-21/06/08 & ESO-VLT-U2 & UVES$^{1}$ & 6700-10425$^{2}$ & 0.027-0.041 \\\\ 2 & 01-09-2009-09/01/2010 & ESO-VLT-U2 & UVES$^{1}$ & 5700-7530 \\& 7650-9470  & 0.027-0.041\\\\ 3 & 03-05/05/2010 & 6.5~m Baade-Maguellan & IMACS Short-Camera$^{3}$ & 6550-10000 & 1.98  \\\\ 4 & 15-16/02/2011 & 6.5~m Baade-Maguellan & IMACS Short-Camera$^{3}$ & 4300-10800 & 1.97  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^{1}$ UVES: Ultraviolet and Visual Echelle Spectrograph.\\\\ $^{2}$ Effective range.\\\\ $^{3}$ IMACS: The Inamori Magellan Areal Camera and Spectrograph.\\\\ \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample A: Model fit results for M3-M7.5 objects. Spectroscopic spectral types are from Paper II. The equivalent effective temperatures for the spectral types were taken from Reyl\\'e et al. (2011). $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resulta}} \\begin{flushleft} \\begin{center} \\small \\begin{tabular}{lcccccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & \\multicolumn{3}{c}{$SpT$} & \\multicolumn{5}{c}{NextGen} & \\multicolumn{4}{c}{Semi-empirical model} & \\\\ \\noalign{\\smallskip} Object & $SpT_{Spec}$ & $T_{\\rm eff}SpT$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & best fit$^{1}$ \\\\ &       & (K)         & & (K) & (cm s$^{-2}$) & dex   &   & & (K) & cm s$^{-2}$ & dex &  & \\\\ \\hline \\noalign{\\smallskip} GL876          & M4.0 & 3100 & & 3400 & 5.5 & 0.0 & 2.9 & & - & - & - & - & NextGen \\\\ SIPS0007-2458  & M7.5 & 2550 & & 2700 & 4.5 & -1.0 & 7.0  & & 2700 & 4.5 & 0.0 & 6.2  & s.-e. \\\\ 2MASS0020-2346 & M6.0 & 2800 & & 2800 & 5.5 & -0.5 & 4.8  & & 2800 & 5.0 & 0.0 & 4.56 & s.-e. \\\\ SIPS0027-5401  & M6.5 & 2650 & & 2700 & 4.5 & -1.0 & 3.35 & & 2700 & 4.5 & 0.0 & 2.70 & s.-e. \\\\ SIPS0153-5122  & M6.0 & 2800 & & 2800 & 5.5 & -0.5 & 2.42 & & 2800 & 5.0 & 0.0 & 2.31 & s.-e. \\\\ SIPS0214-3237  & M6.5 & 2650 & & 2800 & 5.0 & -1.0 & 1.44 & & 2800 & 5.0 & 0.0 & 1.2  & s.-e. \\\\ SIPS0235-0711  & M6.0 & 2800 & & 2700 & 5.0 & -1.0 & 0.97 & & 2700 & 4.5 & 0.0 & 0.91 & s.-e. \\\\ 2MAS0334-2130  & M4.5-M6.0 & 3000-2800 & & 2800 & 5.5 & -0.5 & 4.65 & & 2800 & 5.0 & 0.0 & 5.65 & $^{a}$    \\\\ LEHPM4908      & M6.0 & 2800 & & 2800 & 4.5 & -1.0 & 7.02 & & 2900 & 5.0 & 0.0 & 5.70 & s.-e. \\\\ 2MASS2254-3228 & M5.5 & 2850 & & 2700 & 5.0 & -1.0 & 1.37 & & 2700 & 4.5 & 0.0 & 1.25 & s.-e. \\\\ LEHPM6542      & M6.0 & 2800 & & 2800 & 5.0 & -0.5 & 13.5 & & 2700 & 4.5 & 0.0 & 4.03 & s.-e. \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^{1}$ NextGen or semi-empirical (s.-e.) models.\\\\ $^{a}$ both models are equally likely. \\\\ \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample B:  Model fit results for M3-M7.5 objects. The equivalent effective temperature for the spectral types were taken from Reyl\\'e et al. (2011). $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resultb}} \\begin{flushleft} \\begin{center} \\scriptsize \\begin{tabular}{lccccccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & \\multicolumn{4}{c}{$SpT$} & \\multicolumn{5}{c}{NextGen} & \\multicolumn{4}{c}{Semi-empirical model} & \\\\ \\noalign{\\smallskip} Object & $SpT_{Phot}$ & $SpT_{Spec}$ & $T_{\\rm eff}Sp$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & best fit \\\\ &       & (K)   & (K)      & & (K) & (cm s$^{-2}$) & dex   &  & & (K) & cm s$^{-2}$ & dex &  & \\\\ \\hline \\noalign{\\smallskip} 2MASS1146-4754 & M8.5 & M7-8 & $\\sim$2500  & & 2800 & 5.5 &-0.5 & 0.39 & & 2700 & 4.5 & 0.0 & 0.18 & s.-e. \\\\ 2MASS1236-6536 & M7.5 & M4.0 & 3100       & & 3200 & 5.5 & 0.0 & 0.64 & & 2900 & 4.0 & 0.0 & 1.14 & NextGen \\\\ 2MASS1326-5022 & M9.0 & M7.5 & 2550       & & 2800 & 5.0 & 0.0 & 0.57 & & 2700 & 4.0 & 0.0 & 0.22 & s.-e. \\\\ 2MASS1433-5148 & M9.0 & M6-7 & $\\sim$2650 & & 3000 & 5.5 & 0.0 & 0.58 & & 2900 & 5.0 & 0.0 & 0.27 & s.-e. \\\\ 2MASS1618-3214 & M7.0 & M6.5 & 2650       & & 2800 & 5.5 &-0.5 & 0.69 & & 2700 & 4.5 & 0.0 & 0.41 & s.-e. \\\\ 2MASS1736-0407 & M8.5 & M4-5 & $\\sim$3000 & & 3200 & 5.0 & 0.0 & 0.49 & & 2900 & 4.0 & 0.0 & 0.55 & NextGen  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample A: Model fit results for $\\sim$M8 object. First line refers to fit results when use NextGen, DUSTY and COND models while second line refers to fit results when use the semi-empirical models based on NextGen, DUSTY and COND models respectively. Spectroscopic spectral types are from Paper II. $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resultab}} \\begin{flushleft} \\begin{center} \\scriptsize \\begin{tabular}{lccccccccccccccc} \\noalign{\\smallskip}  \\hline \\noalign{\\smallskip} & \\multicolumn{2}{c}{$SpT$} & \\multicolumn{4}{c}{NextGen} & \\multicolumn{4}{c}{DUSTY} & \\multicolumn{4}{c}{COND} & \\\\ Object & $SpT_{Spec}$ & $T_{\\rm eff}SpT$ &  & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & best fit \\\\ &    &          (K)     & & (K) & (cm s$^{-2}$) & & &  (K) & cm s$^{-2}$ & & & (K) & cm s$^{-2}$ & & \\\\  \\hline \\noalign{\\smallskip} SIPS2039-1126 & M8.0 & ~2500 & & 2800 & 5.5 & 22.45 & & 2600 & 5.0 & 39.96 & & 2600 & 4.5 & 40.12 & \\\\ Semi-empirical & & & & 2700 & 4.5 & 22.90 & & 2600 & 4.5 & 20.00 & & 2600 & 4.5 & 20.01 & s.-e. DUSTY \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample B: Model fit results for $\\sim$M8 objects. First line refers to fit results when use NextGen, DUSTY and COND models while second line refers to fit results when use the semi-empirical models based on NextGen, DUSTY and COND models respectively. $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resultbb}} \\begin{flushleft} \\begin{center} \\scriptsize \\begin{tabular}{lccccccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & \\multicolumn{2}{c}{$SpT$} & \\multicolumn{4}{c}{NextGen} & \\multicolumn{4}{c}{DUSTY} & \\multicolumn{4}{c}{COND} & \\\\ Object & $SpT_{Phot}$ & $SpT_{Spec}$ & & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & $S_{min}$ & best fit \\\\ &    &               & & (K) & (cm s$^{-2}$) & & & (K) & cm s$^{-2}$ & & & (K) & cm s$^{-2}$ & & \\\\  \\hline \\noalign{\\smallskip} 2MASS0814-4020 & L4.0 & M7-8 & & 2700 & 5.0 & 0.74 & & 2600 & 5.5 & 1.15 & & 2600 & 5.0 & 1.22 & \\\\ Semi-empirical & & & & 2700 & 4.5 & 0.35 & & 2600 & 4.5 & 0.28 & & 2600 & 4.5 & 0.28 & s.-e. DUSTY $^{1,2}$\\\\  2MASS1557-4350 & L0.0 & M7.5-8.5 & & 2800 & 4.5 & 0.74 & & 2600 & 5.5 & 0.93 & & 2400 & 5.5 & 1.11 \\\\ Semi-empirical & & & & 2800 & 4.0 & 0.31 & & 2500 & 4.0 & 0.29 & & 2600 & 4.0 & 0.29 & s.-e. DUSTY $^{1}$\\\\  2MASS1756-4518 & M7.5 & M8-9 & & 2600 & 5.5 & 0.21 & & 2600 & 5.5 & 0.18 & & 2600 & 5.0 & 0.20 &  \\\\ Semi-empirical & & & & 2600 & 4.0 & 0.09 & & 2600 & 5.5 & 0.08 & & 2600 & 5.0 & 0.10 & s.-e. DUSTY \\\\  \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^{1}$ DUSTY and COND have the same $S_{min}$, but DUSTY model is more appropriate in the temperature range.\\\\ $^{2}$ We could appreciate parts of the spectrum were NextGen was the best fit.\\\\ \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample A: Model fit results for M8.5-M9.5 objects. First line refers to fit results when use DUSTY and COND models while second line refers to fit results when use the semi-empirical models based on DUSTY and COND models respectively. Spectroscopic spectral types are from Paper II. The equivalent effective temperature for the spectral types were taken from Dahn et al. (2002). $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resultac}} \\begin{flushleft} \\begin{center} \\small \\begin{tabular}{lcccccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & \\multicolumn{3}{c}{$SpT$} & \\multicolumn{5}{c}{DUSTY} & \\multicolumn{4}{c}{COND} & \\\\ \\noalign{\\smallskip} Object & $SpT_{Spec}$ & $T_{\\rm eff}SpT$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & best fit \\\\ &       & (K)         & & (K) & (cm s$^{-2}$) & dex   &   & & (K) & cm s$^{-2}$ & dex &  & \\\\ \\hline \\noalign{\\smallskip} DENIS0021-4244 & M9.5 & 2300-2400 & & 2000 & 5.5 & 0.0 & 7.85 & & 2400 & 5.5 & 0.0 & 13.40 & \\\\ Semi-empirical &  &  & & 2400 & 4.0 & 0.0 & 5.89 & & 2500 & 4.0 & 0.0 & 6.62 & s.-e. DUSTY \\\\ SIPS2045-6332  & M8.5 & 2300-2400 & & 2000 & 5.5 & 0.0 & 11.39 & & 2400 & 4.0 & 0.0 & 17.19 &  \\\\ Semi-empirical &  & & & 2400 & 4.0 & 0.0 & 8.14 & & 2500 & 4.0 & 0.0 & 10.30 & s.-e. DUSTY \\\\ LP944-20       & M9.0 & 2300-2400 & & 2000 & 5.5 & 0.0 & 0.42 & & 2400 & 5.5 & 0.0 & 0.57 &  \\\\ Semi-empirical & & & & 2400 & 4.0 & 0.0 & 0.29 & & 2400 & 4.0 & 0.0 & 0.35 & s.-e. DUSTY \\\\   \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}  \\begin{table*} \\caption[]{Sample B: Model fit results for $>$M8.5 objects. First line refers to fit results when use DUSTY and COND models while second line refers to fit results when use the semi-empirical models based on DUSTY and COND models respectively. The equivalent effective temperature for the spectral types were taken from Reyl\\'e et al. (2011). $\\Delta$T$_{\\rm eff}$= 100 K, $\\Delta$$log~g$= 0.5 cm s$^{-2}$. \\label{tab:resultbc}} \\begin{flushleft} \\begin{center} \\tiny \\begin{tabular}{lccccccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} & \\multicolumn{3}{c}{$SpT$} & \\multicolumn{5}{c}{DUSTY} & \\multicolumn{4}{c}{COND} & \\\\ Object & $SpT_{Phot}$ & $SpT_{Spec}$ & $T_{\\rm eff}SpT$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & & $T_{\\rm eff}$ & $log~g$ & [M/H] & $S_{min}$ & best fit \\\\ &    &       & (K)         & & (K) & (cm s$^{-2}$) & dex   &   & & (K) & cm s$^{-2}$ & dex &  & \\\\ \\hline \\noalign{\\smallskip} 2MASS1734-1151 & L0.0  & M9.0  & 2300-2500 & & 2600 & 5.5 & 0.0 & 1.34 &  & 2600 & 5.0 & 0.0 & 1.43 &  \\\\ Semi-empirical & & & & & 2400 & 4.0 & 0.0 & 0.79 & & 2400 & 4.0 & 0.0 & 0.86 & s.-e. DUSTY \\\\ 2MASS1745-1640 & L8.0  & M9-L2 & $\\sim$2300-2000 & & 2000 & 5.5 & 0.0 & 0.06 & & 2200 & 5.5 & 0.0 & 0.07 &  \\\\ Semi-empirical & & & & & 2000 & 5.5 & 0.0 & 0.05 & & 2200 & 5.5 & 0.0 & 0.05 & s.-e. DUSTY \\\\ 2MASS1909-1937 & M9.0  & L0.0  & $\\sim$2200 & & 2100 & 3.5 & 0.0 & 0.55 & & - & - & - & - &  \\\\ Semi-empirical & & & & & 2000/2100 & 6.0/3.5 & 0.0 & 0.45 & & - & - & - & - & s.-e. DUSTY ? \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}   \\begin{table*} \\centering \\caption[]{MG membership parameters. For each target, we give the result obtained for each criterion used. N/A= non applicable criterion in this case; Y= the age parameter agrees with the MG membership; N= the age parameter does not agree with the MG membership; when more than one membership was possible in column 5, the MG name that the criteria determined as possible is given. Where HY is Hyades, SI is Ursa Major, IC is IC 2391, CA is Castor and PL is Pleiades MG and YD is other young disk object. When interrogation appears after MG name we indicate that the membership probability is fewer than for other MGs or in the case of 2MASS1734-1151 and 2MASS1909-1937, that the criteria are not conclusive.\\\\ \\label{tab:resultfinal}} \\begin{flushleft} \\begin{center} \\scriptsize \\begin{tabular}{lcccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Object & $EW$(Na~{\\sc i})$^{1}$ & $EW$(H$\\alpha$)$^{2}$ & age$^{2}$ & MG memb. & M. from & M. from  & M. from & M. from & M. from & Final MG \\\\ & (\\AA) & (\\AA) & (Myr) & (Kinem/Astrom)$^{3}$ & $EW$(Na~{\\sc i})$^{1}$ & $EW$(H$\\alpha$)$^{2}$ & $log~g$$^{4}$ & Li~{\\sc i}$^{5}$ & $v\\sin{i}$$^{6}$ & \\\\ \\hline \\noalign{\\smallskip} SIPS0007-2458  & 8.2 & 11.0   & 1-300     & IC             & N  & Y  & Y  & N  & Y & {\\it YD} \\\\ 2MASS0020-2346 & 8.1 &  6.1   & $\\sim$300 & YD            & Y  & Y  & Y  & ?  & Y & {\\it YD} \\\\ DENIS0021-4244 & 8.4 & $^{a}$ & N/A        & IC, CA         & CA & N/A & Y  & -  & N & {\\it CA} \\\\ SIPS0027-5401  & 7.7 &  8.1   & $\\sim$300 & HY             & N/A & Y  & Y  & Y  & Y & {\\it HY} \\\\ SIPS0153-5122  & 8.0 &  7.5   & $\\sim$300 & IC, HY         & HY & HY & HY & HY & Y & {\\it HY} \\\\ SIPS0214-3237  & 7.9 &  8.4   & $\\sim$300 & HY             & N/A & Y  & Y  & Y  & Y & {\\it HY}\\\\ SIPS0235-0711  & 7.9 &  7.7   & $\\geq$300 & HY             & N/A & Y  & Y  & Y  & Y & {\\it HY} \\\\ 2MASS0334-2130 & 7.3 & 11.5   & 1-300     & IC             & N  & Y  & N  & N  & ? & {\\it YD }\\\\ SIPS2045-6332  & 6.3 &  2.1   & N/A        & PL, CA         & CA & N/A & Y  & Y  & ? & {\\it CA }\\\\ SIPS2039-1126  & 7.3 &   -    & -         & PL             & N  & -  & Y  & Y  & Y & {\\it PL}\\\\ LEHPM4908      & 7.3 &  6.5   & $\\sim$300 & HY             & N/A & Y  & Y  & Y  & - & {\\it HY }\\\\ 2MASS2254-3228 & 8.5 &  5.8   & $\\sim$300 & PL, CA         & CA & CA & Y  & -  & Y & {\\it CA }\\\\ LEHPM6542      & 7.6 &  4.9   & $\\geq$300 & HY             & N/A & Y  & Y  & Y  & Y & {\\it HY }\\\\ & & &                &    &    & & & & \\\\ 2MASS0814-4020 & 9.0 &  7.2    & $\\geq$300 & HY              & N/A     & Y      & Y      & N/A & N/A & {\\it HY}\\\\ 2MASS1146-4754 & 7.9 &  5.2    & $\\geq$300 & SI              & N/A     & Y      & Y      & N/A & N/A & {\\it SI}\\\\ 2MASS1236-6536 & 5.8 &  1.0    & 300-1200  & PL, IC, CA, SI  & CA, SI & CA, SI & N      & N/A & N/A & {\\it CA, SI} \\\\ 2MASS1326-5022 & 4.2 & 23.6    & 1-100     & IC, CA          & IC     & IC     & Y      & N/A & N/A & {\\it IC}\\\\ 2MASS1433-5148 & 7.9 & 19.8    & 1-100     & PL, IC, CA, SI  & CA, SI & PL, IC & SI     & N/A & N/A & {\\it CA?, SI} \\\\ 2MASS1557-4350 & 4.0 & 16.4    & 1-100     & PL, IC, CA      & PL, IC & N/A     & Y      & N/A & N/A & {\\it PL, IC} \\\\ 2MASS1618-3214 & 8.0 &  7.8    & 1-300     & IC, CA, SI      & CA, SI & Y      & Y      & N/A & N/A & {\\it IC?, CA, SI}\\\\ 2MASS1734-1151 & 7.7 &  5.5    & N/A        & PL              & N      & N/A     & Y      & N/A & N/A & {\\it PL?}\\\\ 2MASS1736-0407 & 6.1$^{c}$ &  4.2    & 50-300    & CA, SI          & N/A     & Y      & Y      & N/A & N/A & {\\it CA, SI}\\\\ 2MASS1745-1640 & 5.6 &  1.4    & $\\geq$300  & CA, SI, HY      & Y      & N/A     & HY     & N/A & N/A & {\\it CA?, SI?, HY} \\\\ 2MASS1756-4518 & 7.5 &  3.2    & N/A        & CA              & N/A     & N/A     & Y      & N/A & N/A & {\\it CA}\\\\ 2MASS1909-1937 & 6.3 & $^{a}$  & N/A        & CA              & N/A     & N/A     & -      & N/A & N/A & {\\it CA?}\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} $^{1}$ Equivalenth width of the Na~{\\sc i} doublet as expained in Section 5.1.\\\\ $^{2}$ See Section 5.2.\\\\ $^{3}$ MG membership from kinematic (sample A) or from photometric and astrometric criteria (sample B). See Section 2.\\\\ $^{4}$ From value obtained in the synthetic fits, given in columns 5 and 9 of Tables~\\ref{tab:resulta}, ~\\ref{tab:resultb} and ~\\ref{tab:resultab}.\\\\ $^{5}$ Only applicable for high resolution data. See Section 5.3.\\\\ $^{6}$ From Paper II.\\\\ $^{a}$ Absorption line filled in with emission.\\\\ $^{b}$ Absorption line.\\\\ $^{c}$ Measured with a cosmic ray in the line.\\\\  \\end{flushleft} \\end{table*}"
    },
    "1402/1402.1967_arXiv.txt": {
        "abstract": "We use a dynamical systems approach based on the method of orthonormal frames to study the dynamics of a two-fluid, non-tilted Bianchi Type I cosmological model. Such a universe is anisotropic, spatially homogeneous, and spatially flat. In our model, one of the fluids is a fluid with bulk viscosity, while the other fluid assumes the role of a cosmological constant and represents nonnegative vacuum energy.  We begin by completing a detailed fixed-point analysis of the system which gives information about the local sinks, sources and saddles. We then proceed to analyze the global features of the dynamical system by using topological methods such as finding Lyapunov and Chetaev functions, and finding the $\\alpha$- and $\\omega$-limit sets using the LaSalle invariance principle. The fixed points found were a flat Friedmann-LeMa\\^{\\i}tre-Robertson-Walker (FLRW) universe with no vacuum energy, a de Sitter universe, a flat FLRW universe with both vacuum and non-vacuum energy, and a Kasner quarter-circle universe. We also show in this paper that the vacuum energy we observe in our present-day universe could actually be a result of the bulk viscosity of the ordinary matter in the universe, and proceed to calculate feasible values of the bulk viscous coefficient based on observations reported in the Planck data. We conclude the paper with some numerical experiments that shed further light on the global dynamics of the system. ",
        "introduction": "In this paper, we use a dynamical systems approach to investigate in detail the dynamics of a Bianchi Type I universe with a bulk viscous fluid and cosmological constant. Such a universe is spatially flat, spatially homogeneous, and anisotropic. Such a model may have considerable importance in present studies of cosmology given the recent results of the Planck measurements \\cite{planckdata}, which suggest that the curvature of the spatial sections of the present-day universe is in agreement with spatial flatness. Moreover, Bianchi models are more general than the Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) models and therefore can provide better descriptions of the early universe where viscous effects may have been dominant \\cite{hervik}. One can then study the effects of viscosity on the dynamical evolution of the universe which we observe to be of FLRW-type today. As discussed by Coley and Wainwright \\cite{coleywainwright}, cosmological models with single fluids are necessarily a simplification in the sense that they can only describe one epoch during the evolution of the universe. More general models can be constructed using two fluids with barotropic equations of state that are also comoving. One can then use these models to describe the transitions between different epochs in the universe's evolution, such as going from a radiation-dominated phase $(w = -1/3)$ to a matter-dominated phase $(w = 0)$. As discussed by Gr{\\o}n and Hervik (Chapter 13, \\cite{hervik}), viscous models have become of general interest in early-universe cosmologies largely in two contexts.  Firstly, in models where bulk viscous terms dominate over shear terms, the universe expands to a de Sitter-like state, which is a spatially flat universe neglecting ordinary matter, and including only a cosmological constant. Such models isotropize indirectly through the massive expansion. Secondly, in the absence of any significant heat flux, shear viscosity is found to play an important role in models of the universe at its early stages. In particular, neutrino viscosity is considered to be one of the most important factors in the isotropization of our universe. By also including a nonnegative cosmological constant in our model, and interpreting it to represent vacuum energy, we are also able to give a detailed description of the roles played by both viscosity and vacuum energy in the isotropization of our universe.  Bianchi cosmological models which contain a viscous fluid matter source in addition to a cosmological constant have been studied in detail several times. Lorenz-Petzold \\cite{lorenzp} examined Bianchi Type I and V models in the the presence of perfect fluid matter with bulk viscosity and a nonzero cosmological constant.  Pradhan and Pandey \\cite{pradhanpandey} studied Bianchi Type I magnetized cosmological models in the presence of a bulk viscous fluid in addition to a monotonically decreasing cosmological constant. Saha \\cite{sahaI} studied the evolution of a Bianchi Type I universe with a viscous fluid and a cosmological constant. Pradhan, Srivastav, and Yadav \\cite{pradhan2} studied Bianchi Type IX viscous models with a time-dependent positive cosmological constant. Belinch\\'on \\cite{belinchon} investigated the dynamics of a locally rotationally symmetric (LRS) Bianchi Type I universe with a bulk viscous fluid and a time-dependent cosmological constant. Pradhan, Jotania, and Rai \\cite{pradhan3} studied Bianchi Type V cosmological models with bulk viscous fluid and a time-dependent cosmological constant. They also discussed some physical and geometrical aspects of such models. Pradhan and Pandey \\cite{pradhan4} studied Bianchi Type I cosmological models with both shear and bulk viscosity and a monotonically decreasing cosmological constant. The authors considered the special case in which the expansion tensor only had two components. Saha and Rikhvitsky \\cite{saha} analyzed a Bianchi Type I universe with a cosmological constant and dissipative processes due to viscosity. They showed that a positive cosmological constant leads to an ever-expanding universe. Singh and Kale \\cite{singhkale2} studied Bianchi Type I, Kantowski-Sachs, and Bianchi Type III cosmological models containing as matter sources a bulk viscous fluid and non-constant gravitational and cosmological constants. Pradhan and Kumhar \\cite{pradhankumhar} studied LRS Bianchi Type II models with bulk viscous fluid and a decaying cosmological constant. Mostafapoor and Gr\\o n \\cite{mostafgron} studied a Bianchi Type I universe with a cosmological constant and nonlinear viscous fluid. Sadeghi, Amani, and Tahmasbi \\cite{sadeghi} investigated a Bianchi Type-VI cosmological model with a cosmological constant and viscous fluid. Barrow \\cite{barrownuc2} showed that models of an inflationary universe driven by Witten strings in the very early universe are equivalent to the addition of bulk viscosity to perfect fluid cosmological models with zero curvature. In this work, Barrow considered the case where the bulk viscosity has a power-law dependence upon the matter density. It was shown that if the exponent is greater than $1/2$, there exist deflationary solutions which begin in a de Sitter state and evolve away from it asymptotically in the future. On the other hand, if this exponent is less than $1/2$, then solutions expand from an initial singularity towards a de Sitter state. Barrow  \\cite{barrownuc3} also estimated the entropy production associated with anisotropy damping in the early universe by considering a Bianchi type I metric with an equilibrium radiation gas and anisotropic stresses produced by shear viscosity. It was shown that the shear viscosity based on kinetic theory has the general form of being proportional to the matter density and that the entropy production due to collisional transport is negligible in such a model.  All of the aforementioned papers use the metric approach (Page 39, \\cite{ellis}) to obtain the dynamical evolution of the Bianchi model under consideration. The alternative approach which is based on the method of orthonormal frames pioneered by Ellis and MacCallum \\cite{ellismac} in conjunction with dynamical systems theory is the path we take in this paper. Belinkskii and Khalatnikov \\cite{belkhat} used phase-plane techniques to study a Bianchi Type I model under the influence of both shear and bulk viscosity.  Goliath and Ellis \\cite{goliathellis} used dynamical systems methods to study FLRW, Bianchi Type I,  Bianchi Type II, and Kantoswki-Sachs models with a positive cosmological constant. Coley and van den Hoogen \\cite{coleyvanV} analyzed in detail a Bianchi Type V model with viscosity, heat conduction, and a cosmological constant. They showed that all models that satisfy the weak energy condition isotropize. Coley, van den Hoogen, and Maartens \\cite{coley} examined the full Israel-Stewart theory of bulk viscosity applied to dissipative FLRW models. Coley and Dunn \\cite{coleydunn} used dynamical systems methods to study the evolution of a Bianchi Type V model with both shear and bulk viscosity. Burd and Coley \\cite{burdcoley} examined using dynamical systems methods the effects of both bulk and shear viscosities upon the FLRW, Bianchi Type I, Bianchi Type V, and Kantowski-Sachs models. They found that these models were structurally stable under the introduction of bulk viscosity. Kohli and Haslam \\cite{isk1} used dynamical systems methods to study the future asymptotic behavior of a Bianchi Type IV model containing both bulk and shear viscosity. Kohli and Haslam \\cite{isk2} used dynamical systems methods to study a Bianchi Type I model containing bulk and shear viscosity in addition to a homogeneous magnetic field.  With respect to dynamical systems methods in multi-fluid models, Stabell and Refsdal \\cite{stabell} considered the dynamics of a two-fluid FLRW system consisting of dust and a cosmological constant. Phase plane methods were used by Madsen, Mimoso, Butcher, and Ellis \\cite{madsenmimoso} to study the evolution of FLRW models in the presence of an arbitrary mixture of perfect fluids. Coley and Wainwright \\cite{coleywainwright} examined orthogonal Bianchi and FLRW models in the presence of a two-fluid system. Ehlers and Rindler \\cite{ehlersrind} studied in great detail three-fluid models, containing radiation, dust, and a cosmological constant. Recently, Barrow and Yamamoto\\cite{barrowyama} considered a two-fluid system with one of the fluids representing a cosmological constant in their study of the instabilities of Bianchi Type IX Einstein static universes. For more details on the history of multi-fluid models, the interested reader should see Pages 53-55, 60-62, 171-172 and references therein of \\cite{ellis}.  Despite all of the important aforementioned contributions, we feel it will be of considerable value to consider the dynamics of a Bianchi Type I universe with a viscous fluid and cosmological constant with respect to dynamical systems theory following the methods outlined in \\cite{ellis} and \\cite{barrowyama}. To the best of the authors' knowledge at the time of writing this paper, such an investigation has not been carried out in the literature.  Throughout this paper, we assume a metric signature of $(-,+,+,+)$ and use geometrized units, where $8\\pi G = c = 1$. ",
        "conclusions": "We have presented in this paper a comprehensive analysis of the dynamical behavior of a Bianchi Type I two-fluid model with bulk viscosity and a cosmological constant. We began by completing a detailed fixed-point analysis of the system which gave information about the local sinks, sources and saddles. We then proceeded to analyze the global features of the dynamical system by using topological methods such as finding Lyapunov and Chetaev functions, and finding the $\\alpha$- and $\\omega$-limit sets using the LaSalle invariance principle.  The fixed points found were a flat FLRW universe with no vacuum energy and only energy due to ordinary matter, a de Sitter universe, a mixed FLRW universe with both vacuum and non-vacuum energy, and a Kasner quarter-circle universe. We found conditions for which the former three were local sinks of the system, that is, future asymptotic states, and where the latter was a source of the system, that is, a past asymptotic state.  The flat FLRW universe solution we found with both vacuum and non-vacuum energy is clearly of primary importance with respect to modelling the present-day universe, especially in light of the recently-released Planck data. In fact, using this Planck data we gave possible conditions for which a non-zero bulk viscosity in the early universe could have lead to some of the conditions described in the Planck data in the present epoch. In particular, since we found that this equilibrium point is a local sink of the dynamical system, all orbits approach this equilibrium point in the future. Therefore, there exists a time period for which our cosmological model will isotropize and be compatible with present-day observations of a high degree of isotropy of the cosmic microwave background in addition to the existence of both vacuum and non-vacuum energy."
    },
    "1402/1402.4129_arXiv.txt": {
        "abstract": "We present the largest search to date for Y-band dropout galaxies ($z\\sim8$ Lyman break galaxies, LBGs) based on 350 arcmin$^2$ of HST observations in the V-, Y-, J- and H-bands from the Brightest of Reionizing Galaxies (BoRG) survey.  In addition to previously published data, the BoRG13 dataset presented here includes approximately 50 arcmin$^2$ of new data and deeper observations of two previous BoRG pointings, from which we present 9 new $z\\sim8$ LBG candidates, bringing the total number of BoRG Y-band dropouts to 38 with $25.5\\leqslant m_{J} \\leqslant 27.6$ (AB system).  We introduce a new Bayesian formalism for estimating the galaxy luminosity function, which does not require binning (and thus smearing) of the data and includes a likelihood based on the formally correct binomial distribution as opposed to the often used approximate Poisson distribution. We demonstrate the utility of the new method on a sample of $97$ Y-band dropouts that combines the bright BoRG galaxies with the fainter sources published in \\cite{Bouwens:2011p8082} from the Hubble Ultradeep Field (HUDF) and Early Release Science (ERS) programs.  We show that the $z\\sim8$ luminosity function is well described by a Schechter function over its full dynamic range with a characteristic magnitude $M^\\star = -20.15^{+0.29}_{-0.38}$, a faint-end slope of $\\alpha = -1.87^{+0.26}_{-0.26}$, and a number density of $\\log_{10} \\phi^\\star [\\textrm{Mpc}^{-3}] = -3.24^{+0.25}_{-0.24}$.  Integrated down to $M=-17.7$ this luminosity function yields a luminosity density, $\\log_{10} \\epsilon [\\textrm{erg}/\\textrm{s/Hz/Mpc}^{3}] = 25.52^{+0.05}_{-0.05}$. Our luminosity function analysis is consistent with previously published determinations within 1$\\sigma$. The error analysis suggests that uncertainties on the faint-end slope are still too large to draw firm conclusion about its evolution with redshift.  We use our statistical framework to discuss the implication of our study for the physics of reionization.  By assuming theoretically motivated priors on the clumping factor and the photon escape fraction we show that the UV luminosity density from galaxy samples down to $M=-17.7$ can ionize only 10-50\\% of the neutral hydrogen at $z\\sim8$. Full reionization would require extending the luminosity function down to $M=-15$. The data are consistent with a substantial fraction of neutral hydrogen at $z>7$, in agreement with recent suggestions based on deep spectroscopy of redshift 8 LBGs. ",
        "introduction": "\\label{sec:intro}  Characterizing the epoch of reionization, i.e., the epoch at which the first stars and galaxies in the Universe reionized the vast majority of the neutral hydrogen, is a key outstanding issue in the continued effort to map the formation and early evolution of galaxies. For almost four years, since the installment of the Wide Field Camera 3 (WFC3) onboard the Hubble Space Telescope (HST), the frontier of this knowledge has been pushed further and further back towards the dawn of cosmic reionization.  At present, several hundred high redshift galaxy candidates have been found at redshift $z\\sim6$ \\citep[e.g.,][]{Stark:2010p27668,Bouwens:2007p29848,Bouwens:2012p31891,Ouchi:2010p27936,Bradley:2013p32053}, and with the improved near-IR efficiency of WFC3 the search for candidates has been pushed to $z\\gtrsim 8$ using the Lyman Break technique. In particular the Hubble Ultra Deep field efforts in 2009 \\citep{Oesch:2010p30140,Oesch:2010p30134,Lorenzoni:2011p12992,Bouwens:2010p30142,Bouwens:2011p8082,McLure:2010p31124} and 2012 \\citep{Koekemoer:2012p26719,Ellis:2012p26700,Dunlop:2012p26717,Schenker:2013p26914,McLure:2013p27183,Ono:2012p26883,Oesch:2013p27877,Illingworth:2013p31165} have revealed the faintest samples of galaxy candidates at $z\\gtrsim8$. Simultaneously larger area observations are targeting brighter and rarer candidates, either in legacy fields, such as GOODS/CANDELS \\citep{Grogin:2011p9457,Koekemoer:2011p9456}, or in pure-parallel random pointings, like those of our \\emph{Brightest of Reionizing Galaxies Survey}\\footnote{\\url{https://wolf359.colorado.edu}} \\citep[hereafter BoRG,][]{Trenti:2011p12656,Trenti:2012p13020,Bradley:2012p23263}.  Our ongoing BoRG survey has two key goals. The first goal is to provide bright targets that can potentially yield spectroscopic confirmation of $z\\sim8$ galaxies by follow-up observations \\citep{Treu:2012p12658,Treu:2013p32132}. In fact, while $z\\sim6$ dropout samples have extensive spectroscopic redshifts, \\citep[e.g.,][]{Vanzella:2009p29479,Stark:2010p27668,Stark:2013p28962}, only a handful of $z\\gtrsim 7$ galaxies \\citep[e.g.,][]{Ono:2012p27651} have currently confirmed redshifts with the highest being at $z=7.5$ \\citep{Finkelstein:2013p32467}. So far no Y-band dropout at $z\\sim8$ has been spectroscopically confirmed. Only upper limits on Ly$\\alpha$ flux have been provided to date \\citep{Caruana:2012p27502,Caruana:2013p32713,Capak:2013p31627,Treu:2013p32132,Faisst:2014p34184} and those leave open the interpretation of whether the photometric selection technique breaks down at $z\\gtrsim 7$ (which would be surprising given the small change in magnitudes and filters) versus the more interesting physical explanation of an increase in the intergalactic medium (IGM) optical depth to Ly$\\alpha$ arising from a higher neutral hydrogen fraction at $z\\sim8$ \\citep{Treu:2013p32132} with respect to  redshift 7 \\citep{Fontana:2010p29506} and 6 \\citep{Stark:2010p27668,Stark:2013p28962}.  The second goal of the BoRG survey is to improve the determination of the $z\\sim8$ luminosity function, by identifying rare and bright dropouts to extend the dynamic range of observations in smaller area deep fields, which are dominated by fainter sources. An accurate measure of the luminosity function is necessary not only to study how galaxies evolve across time, but also to quantify the photon budget available for hydrogen ionization \\citep{Trenti:2010p29335,Zaroubi:2013p24088,Dunlop:2013p23759}. At lower redshift, it is well established \\citep{Bouwens:2007p29848} that the luminosity function is accurately described by a Schechter function \\citep{Schechter:1976p29330} so it is natural to expect a similar form at higher redshift. However, data covering a wide dynamic range are needed to establish that this is indeed the case, and to resolve the degeneracy between the Schechter function parameters in the luminosity function fit \\citep{Bradley:2012p23263,Oesch:2012p30149}.  In this paper we have two goals. The first is to present the complete sample of Y-band dropouts from the BoRG cycle-19 data, and to use these in combination with the literature to determine the galaxy luminosity function at $z\\sim8$. The second goal is to study the consequences for cosmic reionization the inferred luminosity function has. To determine the luminosity function we develop and implement a rigorous statistical Bayesian method to infer the posterior distribution function of the parameters of the luminosity function from the data. The method supersedes those commonly adopted in this field \\cite[e.g][]{Bradley:2012p23263,McLure:2013p27183,Schenker:2013p26914,Oesch:2012p30149} in several ways: the data are not binned, thus avoiding smearing the luminosity function \\citep{Trenti:2008p32309}; the flux uncertainties are correctly taken into account; the counts are modeled using the formally correct binomial distribution \\citep{Kelly:2008p29070}, instead of the Poisson approximation; the full posterior probability distribution function is computed using Markov Chain Monte Carlo methods instead of relying on maximum likelihood estimators based on the asymptotic covariance matrix from the observed Fisher information for uncertainties. By applying this framework to a large sample of $z\\sim 8$ galaxies, consisting of $N=97$ objects both bright (from BoRG) and faint (from the Hubble UDF/ERS fields), we show that the credible intervals include previous best-fit estimates By treating the problem in a self-consistent statistical manner we carry out an inference about reionization by combining the inferred observational uncertainties with various theoretical priors.  The paper is organized as follows. We start by briefly describing the BoRG survey in Section~\\ref{sec:borg}. The current sample of Y-band dropouts containing 9 new and 2 improved $z\\sim8$ galaxy candidates (BoRG13) with respect to those previously published by our team \\citep[BoRG09; BoRG12][]{Trenti:2011p12656,Bradley:2012p23263} is described in Section~\\ref{sec:ydrop}. Two Appendixes (\\ref{app:borg58} and \\ref{app:noise}) take advantage of the follow-up observations of one field and of the large number of BoRG pointings to characterize and discuss the statistics of detections and contaminants in dropout searches. In Section~\\ref{sec:LF} we apply our inference of the Schechter luminosity function parameters using our rigorous Bayesian framework, discussed in detail in Appendix~\\ref{sec:BF}.  The results are presented and discussed in the context of cosmic reionization in Section~\\ref{sec:results}. A brief summary is given in Section~\\ref{sec:conc}.  All magnitudes are AB magnitudes and a standard concordance cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $h=0.7$ is assumed. ",
        "conclusions": "\\label{sec:conc}  The BoRG survey has carried out the largest-area search to date for Y-band dropouts (HST F098M-dropouts). We present new observations from 12 parallel fields not included in our previous studies and additional deeper datasets for two fields. We combine our BoRG sample with a sample of fainter dropouts taken from the literature and use the combined sample of $97$ dropouts to carry out a rigorous study of the luminosity function at $z\\sim8$ and its implications for reionization. Our main results can be summarized as follows: \\begin{enumerate} \\item We present 9 new Y-band dropouts from the $\\sim$50 arcmin$^2$ of new data. Furthermore, we re-confirm two dropouts previously presented by \\citet{Trenti:2011p12656} and \\citet{Bradley:2012p23263}. Combining these 11 dropouts with out previously published $z\\sim8$ LBGs from BoRG gives a sample of 38 bright ($25.5\\leqslant m_{J} \\leqslant 27.6$) redshift 8 LBGs from the BoRG survey.  \\item We develop and implement an improved method for estimating the luminosity function parameters for a sample of $n$ galaxies. Using a Bayesian framework the posterior distribution of the population based on a binomial distribution is given. Combining the BoRG Y-band dropouts with the faint $z\\sim8$ LBGs from HUDF/ERS presented by \\cite{Bouwens:2011p8082} and sampling over this posterior distribution enables a robust recovery of the faint-end slope, $\\alpha$, the characteristic magnitude of the Schechter function, $M^\\star$, the normalizing number density of $z\\sim8$ LBGs, $\\phi^\\star$, and the luminosity density, $\\epsilon$, of the redshift 8 luminosity function.  \\item The inferred luminosity function at $z\\sim8$ is described by the parameters: \\begin{itemize} \\item[] $M^\\star= -20.15^{+0.29}_{-0.38}$, \\item[] $\\alpha= -1.87^{+0.26}_{-0.26}$, \\item[] $\\log_{10} \\phi^\\star [\\textrm{Mpc}^{-3}]=  -3.24^{+0.25}_{-0.34}$, \\item[] $\\log_{10} \\epsilon \\left[\\frac{\\textrm{erg}}{\\textrm{s Hz Mpc}^{-3}}\\right]= 25.52^{+0.05}_{-0.05}$ \\end{itemize} Here $\\epsilon$ is the inferred UV luminosity density integrated down to the HUDF limit M$_{\\rm lim}=-17.7$. Our inferred credible intervals include recent estimates of the same parameters.  \\item We show that for the BoRG13 $z\\sim8$ luminosity function the average fraction of ionized hydrogen $Q$ is only of the order 10-50\\% for samples down to the HUDF limit of $M=-17.7$ assuming standard values of the clumping factor and the photon escape fraction. To sustain a fully ionized Universe at redshift 8 with the presented luminosity function it is necessary to account for the radiation of objects as faint as $M=-15$.  \\item The inferred ionization fractions suggest a relatively late reionization scenario where a considerable fraction of neutral hydrogen is still present at $z\\sim8$ in good agreement with the results of our recent spectroscopic MOSFIRE campaign (and others from the literature) where we followed up a subsample of the BoRG redshift 8 LBGs presented here. \\end{enumerate}  The inference on the implications of the BoRG13 luminosity function for reionization presented here are still limited by the sizable error bars at redshift 8. To reduce the uncertainties on the inferred quantities bright and faint galaxy samples from e.g., the Frontier Fields and future parallel HST campaigns are crucial."
    },
    "1402/1402.4403_arXiv.txt": {
        "abstract": "While Population III stars are typically thought to be massive, pathways towards lower-mass Pop III stars may exist when the cooling of the gas is particularly enhanced. A possible route is enhanced HD cooling during the merging of dark-matter halos. The mergers can lead to a high ionization degree catalysing the formation of HD molecules and may cool the gas down to the cosmic microwave background (CMB) temperature. In this paper, we investigate the merging of mini-halos with masses of a few 10$^5$~M$_\\odot$ and explore the feasibility of this scenario. We have performed three-dimensional cosmological hydrodynamics calculations with the \\verb|ENZO| code, solving the thermal and chemical evolution of the gas by employing the astrochemistry  package \\verb|KROME|. Our results show that the HD abundance is increased by two orders of magnitude compared to the no-merging case and the halo cools down to $\\sim$60 K triggering fragmentation. Based on Jeans estimates the expected stellar masses are about 10 M$_\\odot$. Our findings show that the merging scenario is a potential pathway for the formation of low-mass stars. ",
        "introduction": "The formation modes of the first generations of stars have been a challenging subject for many years and today still represent a central topic of research in modern cosmology. The current understanding is that these stars were formed in mini-halos of about 10$^6$ M$_\\odot$ around $z$~=~20-30, which is strongly supported by many numerical simulations performed over the years \\citep{Abel2002,Bromm2002,Ciardi05,Yoshida2006,Oshea2007,Greif2008,Greif2012}. Since \\citet{Saslaw1967}, it is widely accepted that the collapse process leading to the formation of these early objects is mainly driven by H$_2$ cooling \\citep{Peebles1968,Palla1983}. Due to the lack of a permanent dipole moment the hydrogen molecule becomes an ineffective coolant below $\\sim$200 K when it reaches the local thermodynamic equilibrium (LTE). Considering these thermal conditions, it has been postulated that the final mass of the first generation of stars should be of the order of $\\sim$100~M$_\\odot$ or more leading to a resultant short life \\citep{Abel2002,Bromm2002,BrommRev2013}. However, recent numerical simulations \\citep{ClarkGlover2011,Greif2012,Stacy2012} suggest that the characteristic mass scale of the first stars should be lowered to several ten solar masses, if fragmentation occurs. The large parameter study by \\citet{Hirano2013} indicates a significant spread around this mass scale, and extreme values up to 1000 M$_\\odot$.  Comprehending the conditions under which the first stars were formed is very important to study the early history of the Universe as they initiated cosmic reionization and the enrichment of the intergalactic medium \\citep[e.g.][]{Barkana2001,Schneider06,Schleicher08}, and may also have seeded the formation of supermassive black holes \\citep{Li2007,Volontieri2012,LatifBH,Hosokawa2013,LatifBH2}.  Many authors \\citep{Bromm2002,Johnson2006MNRAS.366..247J,Yoshida2006,Ripamonti2007,Greif2008,McGreer2008} also discussed the possible role of HD as an additional coolant which may have led to the formation of lower-mass stars. They found that in the presence of a higher electron abundance, HD cooling can become efficient allowing the gas to reach the temperature of the CMB. The need of a high ionization fraction suggested the introduction of a second mode to form primordial stars named Pop~III.2 to be distinguished from Pop~III.1 in which H$_2$ cooling dominates throughout \\citep[see][]{Tan2008,BrommRev2013}. The Pop III.2 stars are expected to be assembled from a gas which is still metal-free but embedded in ionized environments. The latter may be triggered by the explosions of the first supernovae or HII regions or strong shocks \\citep[e.g.][]{Mackey2003,Yoshida2007,Greif2008}. Lower-mass stars may also be formed within the Pop III.1 paradigm as shown by \\citet{Hirano2013}. They suggested that in a slowly collapsing cloud the compressional heating is reduced and may create the conditions for the HD cooling to become effective. In some of the above mentioned papers \\citep{Bromm2002,Ripamonti2007,McGreer2008} the authors claimed that in mini-halos of a few $\\sim$10$^5$~M$_\\odot$ HD can efficiently be formed and overcome the H$_2$ cooling rate if the temperature drops below $\\sim$200~K at densities of at least 10$^4$ cm$^{-3}$. \\citet{Stacy2013} performed cosmological simulations of the formation and growth of Pop III stellar systems and argued that a mini-halo formed at redshift 15 and supported by high baryonic angular momentum can lead to the formation of 1-5 M$_\\odot$ stars even in the absence of HD cooling. Lower-mass stars are thus also included in the expected scatter of the Pop III.1 distribution.  \\citet{McGreer2008} concluded that in halos with masses $>$10$^6$ M$_\\odot$ the cooling is dominated by H$_2$ and the final stars could have a mass of $\\sim$100~M$_\\odot$, while for lower mass halos HD cooling becomes important and leads to stellar masses between 10-40~M$_\\odot$. However, the occurrence of such low-mass systems was found to be less frequent. These conclusions are in overall agreement with the 1D Lagrangian simulation reported by \\citet{Ripamonti2007} and with the robust statistical study presented by \\citet{Hirano2013}.  The main idea behind the possible role of HD in primordial star formation comes from the series of chemical reactions which leads to the formation of HD. This is strongly related to the ionization fraction because the main formation path for HD is \\begin{equation} \\mathrm{H_2 + D^+ \\rightarrow HD + H^+}\\\\, \\end{equation} as largely discussed in earlier papers \\citep{Stancil1998,Galli2002,Stancil2011}. The only way to trigger the formation of HD is then to boost the formation of H$_2$ pumping the main reactive channel which is directly linked to H$^-$ via the following reactions: \\begin{eqnarray} &\\mathrm{H} + \\mathrm{e^-} \\rightarrow \\mathrm{H^-} + \\gamma\\\\ &\\mathrm{H^-} + \\mathrm{H} \\rightarrow \\mathrm{H_2} + \\mathrm{e^-}. \\end{eqnarray} In addition the formation and destruction of D$^+$ should be considered \\begin{equation} \\mathrm{D^+} + \\mathrm{H} \\rightleftharpoons \\mathrm{H^+} + \\mathrm{D}. \\end{equation} This leads us to the conclusion that a high ionization fraction boosts the formation of HD via a chain of chemical events that includes the key species in the following way $\\mathrm{e^-\\rightarrow H^-\\rightarrow H_2 \\rightarrow HD}$. In addition, we note that also an increase of the gas temperature will boost the formation of H$^{-}$, and thus H$_2$ and HD.  Another possible route for primordial star formation (Pop III.2) which also involves the HD cooling has been proposed by \\citet{Merging2006}, and \\citet{Prieto2012,Prieto2013}. The idea suggested by \\citet{Merging2006} is based on the merging of dark-matter halos within the context of the hierarchical scenario of structure formation \\citep{Barkana2001,Ciardi05}. Merging induces compression and the consequent formation of shock-waves which may enrich the baryonic component of the gas by free electrons catalyzing the HD formation. The authors consider a simplified one-dimensional model where two identical halos collide and follow the non-equilibrium evolution of the HD behind the shock-waves which formed during the merging. In these halos, they further assume a uniform density equal to the virial density. They found that to have a higher ionization degree and to boost the formation of HD the merger masses should be $M>8\\times 10^6[(1+z)/20]^{-2}$~M$_\\odot$. In a real merger, we however expect a more complex density structure, and the strength of the shocks is likely enhanced in the central regions. Their mass scale can therefore just serve as a rough estimate. In a recent study, \\citet{Prieto2012,Prieto2013}  performed cosmological N-body simulations and computed the statistic of the halos fulfilling the \\citet{Merging2006} criterion. Their findings suggest that the fraction of halos going through this phase is significant.  In this paper we investigate the role of HD as a coolant in post-shocked environments generated by dark-matter halo merging. In particular we consider a merger of more than two halos with masses well below the threshold set by \\citet{Merging2006} and run self-consistent cosmological simulations, following both the merging and the collapsing process with enough accuracy to capture important chemo-dynamical features.  The paper is organized as follow: we first introduce the methodology and the chemical model involved in the cosmological simulations, then discuss the main results, and finally give our conclusions. ",
        "conclusions": "\\label{sect:conclusions} In this paper we follow the evolution of a mini-halo of $\\sim$7$\\times$10$^5$ M$_\\odot$ formed from the merging of a number of halos with masses of a few 10$^5$ M$_\\odot$. It was in fact suggested \\citep{Merging2006} that if two massive enough halos collide and merge the gas can go through a post-shocked ionized stage which boosts the formation of the HD molecule then allowing the gas to cool below 200 K. As the final mass of the first luminous objects is given by the Jeans mass \\begin{equation} M_J = 500\\,\\mathrm{M_\\odot}\\left(\\frac{T}{200 \\mathrm{K}}\\right)^{3/2}\\left(\\frac{10^4 \\mathrm{cm^{-3}}}{n}\\right)^{1/2} \\end{equation} it is clear that in an environment dominated by H$_2$ cooling the expected mass is of about 500 M$_\\odot$ (assuming $T=200$ K and $n = 10^4$ cm$^{-3}$), while for the HD cooling case the mass is lowered to $\\sim$~10 M$_\\odot$ (e.g. assuming a temperature of 60~K at $n\\sim10^6$ cm$^{-3}$).  A first attempt to investigate the merger-induced HD cooling was made by \\citet{Prieto2012} who performed hydro-cosmological simulations of 3$\\times$10$^7$ M$_\\odot$ halo resulting from the merging of two minihalos of a few 10$^6$ M$_\\odot$. A non-equilibrium treatment of the chemistry was included as well as cooling from H$_2$ and HD. They followed the evolution of the halo until densities of 10$^3$ cm$^{-3}$. The main results from their work was mainly focused on development of turbulence during the merging process. However, they also noticed that few cells of the entire simulations were able to reach temperatures of 100 K with HD fraction around 10$^{-6}$. The densities reached in their work are nevertheless too low to catch possible features due to HD overcooling.  \\citet{Prieto2013} continued their investigation studying the statistics of the halo going through a merging process within a cosmological framework but performing DM-only simulations. They suggested that $\\sim$30\\% of the halos which generates from mergers are able to overcool the gas.  Here we go beyond these studies by coupling the solution of dark matter and hydrodynamics (solved with \\verb|ENZO|) with an accurate treatment of the non-equilibrium chemistry (solved with \\verb|KROME|) following the collapse until a gas density of 10$^8$ cm$^{-3}$, thus extensively covering the density regime where HD cooling is known to usually dominate \\citep{Lipovka2005,GloverAbel08}. We performed high-resolution cosmological simulations for a box of 1 comoving Mpc and follow the evolution of the halo starting from $z=99$, going through the end of the merging process around $z=12$, and following the collapse of our final halo of $\\sim$7$\\times$10$^5$ M$_\\odot$ until $z=11.42$.  Our results clearly show that the shock-waves generated by the merging process enhance H$^-$ and D$^+$ fractions which catalyze the formation of HD. We found a core HD abundance two orders of magnitudes larger than the one obtained with an isolated halo setup. As a consequence, the temperature in the core dropped down to 60-70 K, much lower compared to the isolated halo run where the gas is hotter (around $\\sim$500 K).  Such over-cooling in the halo has induced fragmentation and multiple clumps are observed, out of which two are rotationally supported and one is gravitationally bound.  \\begin{figure*} \\begin{center} \\includegraphics[width=.8\\textwidth]{fig4OK} \\caption{Density, temperature, and HD mass fraction projections, for the run where merging occurs at a scale of 10 pc for two different redshifts, $z=11.44$ (upper panel), and $z=11.43$ (lower panel).}\\label{fig:figure4} \\end{center} \\end{figure*}   It is quite interesting to note that a merging process involving several halos with masses of few 10$^5$~M$_\\odot$ can similarly catalyze the required HD fraction via the efficient production of H$^-$. The differences in the masses of the mergers rise from the fact that we performed realistic three-dimensional calculations compared to the simplified one-dimensional estimates by \\citet{Merging2006} that provide only a qualitative picture of the process. In particular simplified one-dimensional models underestimate the strength of the merger shocks and do not take into account the collapse dynamics and the hydrodynamical effects (e.g. compressional heating). Another strong simplification is the assumption of a virial density while realistic halo are partially collapsed, enhancing the formation of molecules. Based on the halo finder results from the simulations we found a relative velocity of $\\sim$7.0 km s$^{-1}$ between the dark matter halos which is about a factor of 1.5 larger than the estimate by \\citet{Merging2006} from their Eq. 2.  In a previous study, \\cite{McGreer2008} found that the fraction of HD-cooled halos is negligible in primordial \\emph{unperturbed} environments for halos of masses $\\geq$10$^6$~M$_\\odot$, but can become relevant for masses $<$10$^6$~M$_\\odot$ or in ionized media. The authors explored a series of physical conditions that should favour the formation of HD, i.e. gas efficiently cooled by H$_2$ to a critical temperature which should hold for time long enough to allow the build up of HD molecules. These conditions are more suitable in lower mass halos which collapse more uniformly. On the other hand, \\citet{Ripamonti2007} and \\citet{Bromm2002} have shown that the critical masses should be of a few 10$^5$~M$_\\odot$ and attribute the increase in HD fraction to a longer collapsing time. This was recently confirmed by the comprehensive study by \\citet{Hirano2013}. However, in the study by \\cite{McGreer2008}, it is not clear if mergers have occurred during the halo formation then a direct comparison with our results is difficult. Here, we simulate a situation where the halo formed via the occurance of major mergers and we find that HD abundance is enhanced up to 10$^{-5}$ in the aftermath of the merging process. Our findings show that the merger-driven shocks enhance the temperature on large scales, thus favouring the formation of H$^-$ and increasing the H$_2$ fraction. Also the HD abundance is slightly enhanced at that stage, even though it is unable to produce a runaway effect at those densities. When the collapse occurs, the enhanced H$_2$ and HD abundance however allows to trigger a runaway effect as also described by \\citet{McGreer2008}. Our findings suggest that the fraction of these halos might be larger than previously estimated as merger events are much more common in the early Universe ($\\sim$30\\% as reported in \\citet{Prieto2012}).  The present work indicates a potential formation route for low-mass primordial stars due to merger-enhanced HD cooling. This pathway should be explored further in future studies, including a larger range of initial conditions.  Finally, the presence of high turbulent energy produced in the aftermath of merging event may also contribute to the amplification of magnetic fields through the so-called small scale dynamo process \\citep{Schober2012ApJ,Bovino2013NJP,Schleicher2013AN,LatifBH3,Schleicher2013NJPh,Latif2013MNRAS}. The presence of such fields may further influence fragmentation and disk formation at higher densities \\citep[e.g.][]{Machida2013,Latif2013}."
    },
    "1402/1402.3824_arXiv.txt": {
        "abstract": "{This article is the second in a two part series introducing r-Java 2.0, a nucleosynthesis code for open use that performs r-process calculations and provides a suite of other analysis tools.} {The first paper discussed the nuclear physics inherent to r-Java 2.0 and in this article the astrophysics incorporated into the software will be detailed. } {R-Java 2.0 allows the user to specify the density and temperature evolution for an r-process simulation.   Defining how the physical parameters (temperature and density) evolve can effectively simulate the astrophysical conditions for the r-process.  Within r-Java 2.0 the user has the option to select astrophysical environments which have unique sets of input parameters available for the user to adjust.  In this work we study three proposed r-process sites; neutrino-driven winds around a proto-neutron star, ejecta from a neutron star merger and ejecta from a quark nova.  The underlying physics that define the temperature and density evolution for each site is described in this work.} {In this paper a survey of the available parameters for each astrophysical site is undertaken and the effect on final r-process abundance is compared.  The resulting abundances for each site are also compared to solar observations both independently and in concert.  R-Java 2.0 is available for download from the website of the Quark-Nova Project: \\url{quarknova.ucalgary.ca} } {} ",
        "introduction": "\\label{intro}  The rapid neutron capture process (r-process) is believed to be the mechanism for the nucleosynthesis of about half of the stable nuclei heavier than iron \\citep{Burbidge,Cameron}.  Explosive and neutron-rich astrophysical environments present ideal conditions for the r-process to take place. Possible candidate sites discussed in the literature include the neutrino-driven neutron-rich wind from proto-neutron stars \\citep{qian96,Qian03b}, prompt explosions of collapsed stellar cores, \\citep{Sumiyoshi,Wanajo,sar12}, neutron star decompression \\citep{meyer97,Goriely}, tidal disruption in binary merger events \\citep{FRT}, outflows in gamma-ray bursts \\citep{Surman}, the LEPP process in low metalicity stars \\citep{trav04}, supernova fallback \\citep{fryer06}, etc.  Most importantly, abundance data on r-process elements in  metal-poor stars \\citep{Sneden03,Truran} and certain radio-nuclides in meteorites \\citep{QW08} point toward the distinct possibility of multiple r-process sites.  An important paper by \\cite{beth85} on supernova neutrinos brought the high-entropy environment in neutrino-driven winds from Type II Supernovae (SNe) to the forefront of the discussion about astrophysical sites for the r-process.  Since then, much progress has been made in the modelling of type II SNe and neutrino winds of nascent neutron stars \\citep{Woosley,Takahashi,qian96,Cardall,Otsuki,Wanajo01,TBM}. The most natural explanation in the neutrino-driven wind scenario is that the observed r-process pattern follows from a superposition of neutron capture events with differing neutron-to-seed ratios and exposure time-scales. A particular challenge for high-entropy winds as an r-process site is that producing the third peak requires extreme values of entropy and dynamic time-scale that are not supported by current hydrodynamic models of Type II SNe explosions \\citep[e.g.][]{arc07,QW08,fis10, arc11}.  Neutron star mergers can provide a much larger neutron-to-seed ratio than type II SNe, which makes for an appealing r-process site.  Recent relativistic hydrodynamic simulations of neutron star mergers have shown that a significant amount of r-process enriched matter can be ejected \\citep{janka99,ross04,oec07,gori11}, implying that the solar abundance of r-process nuclei may be influenced by neutron star mergers.  However the neutron star (NS) coalescences time scale which likely ranges from 1 - 1000 Myr$^{-1}$ per Milky Way Equivalent Galaxy \\citep{kalo,kaloErr} puts the neutron star merger explanation at odds with enrichment of r-process elements relative to iron observed in metal-poor stars \\citep{qian00, arg04}.  A possible new site for the r-process that is conjectured to be effective in producing the heavy elements beyond the second peak is the quark nova \\citep{ODD}. If this scenario occurs in nature, it presents a new possibility for explaining the origin of the heavy elements. \\cite{meyer89} and \\cite{Goriely} have studied the r-process in decompressing {\\it cold} neutron star matter \\citep{latt77}, although no specific mechanism for decompression was proposed. In previous work \\citep{jaik07,nieb10}, it was suggested that the dynamics of a quark-hadron phase transition inside a neutron star could be sufficiently strong to power the decompression and subsequent ejection of the neutron-rich crust, this is essentially the quark nova scenario.   Clearly there is further study needed to better understand r-process in the Universe and to this end we present r-Java 2.0 an r-process code that is transparent and freely available\\footnote{Software and user manual can be downloaded at quarknova.ucalgary.ca} to the nuclear physics and astrophysics community.  This article is meant to display the astrophysics incorporated in r-Java 2.0 and should be considered in conjunction with \\cite{kostka} which covered the nuclear physics inherent to the software.  This paper is organized as follows; an overview of r-Java 2.0 is presented in section \\ref{rJavaNuc}, the high-entropy wind (HEW) module is described in section \\ref{hewSect}, the neutron star merger (NSM) module is laid out in section \\ref{nsmSect} and the quark nova (QN) module is discussed in section \\ref{qnSect}.  Within each astrophysical module section simulation results are presented.  In section \\ref{discuss} the results from each astrophysical scenario is compared to the solar observations.  Finally a summary is presented in section \\ref{conc}. ",
        "conclusions": "\\label{conc} For the HEW simulations done in this work when the initial nuclei abundances were calculated under the assumption of NSE the r-process was not capable of producing nuclei significantly heavier that A = 80.  When an $\\alpha$-rich freeze-out was assumed to compute the initial abundances a strong A = 130 peak was synthesized for all parameters surveyed.  Starting from an $\\alpha$-rich freeze-out only the largest entropy (S = 270 and thus neutron-to-seed ratio = 100) studied was capable of producing transuranic elements.  The HEW scenario reproduced the europium rare earth peak observed in solar r-process abundances.  Fission recycling had the most impact on final r-process abundances in the NSM of all the scenarios studied.  For all parameters surveyed (with the exception of the initial abundance of only neutrons, protons and alpha particles) r-process in the context NSM was capable of synthesizing the A = 195 peak and in most cases transuranic elements as well.  When compared to the solar abundances NSM robustly reproduced the A = 130 and A = 195 peaks, however the europium rare earth peak was not formed.  Recent work by \\cite{gori13} studying fission fragmentation of neutron-rich nuclei has addressed the europium rare earth peak in the context of NSM.  Simulations done in this work showed that an ejecta mass of $10^{-4} M_{\\odot}$ produced the strongest A =195 of all QN scenarios studied.  The lighter ejecta mass ($10^{-5} M_{\\odot}$) studied cooled too rapidly for the r-process to be able to produce transuranic elements.  The final r-process abundances in heavier ($10^{-3} M_{\\odot}$) ejecta case produced a relatively flat distribution of nuclei.  When compared to solar abundances, the QN scenario under-produced the breadth of the A = 130 peak but was capable of synthesizing the rare earth peak and also the A = 195 peak.  For the simulations performed in this work no single parameter set for a chosen astrophysical site could reproduce the solar r-process abundance distribution.  Either a combination of different astrophysical sites or different parameter sets within an astrophysical site was needed in order to re-produce solar abundances.  Within the parameters surveyed in this work each scenario is capable of producing both a A = 130 and A = 195 peak comparable to those observed in the solar abundances.  For the parameters studied in this work, the resultant r-process abundances for the NSM and QN were quite similar.  This similarity is due to the high initial neutron-to-seed ratio for both scenarios which causes fission recycling to become a driving factor that shapes the final nuclei distribution.  The similarity in r-process yields from these two scenarios could be an indication as to a cause for the apparent universality of the r-process nuclei distribution \\citep{qianUni}.  The study of r-process nucleosynthesis remains a challenging topic which requires a deeper understanding of both the underlying nuclear physics as well as the astrophysics that shapes the r-process environment.  With r-Java 2.0 we have provided a robust and easy-to-use platform that is capable of tackling both sides of the r-process problem.  In \\cite{kostka} we discussed the cutting-edge nuclear physics incorporated in r-Java 2.0 which includes; temperature-dependant neutron capture cross-sections and photo-dissociation rates, realistic fission recycling and beta-delayed neutron emission.  In this work we have shown the capabilities of r-Java 2.0 to model the evolution of density and temperature of three different proposed astrophysical r-process site; HEWs around a proto-neturon star, ejecta from NSMs, and ejecta from QNe.  A fundamental tenet followed during the development of r-Java 2.0 was maximizing flexibility.  For this reason we ensured that the software is cross-platform compatible and gave the user the ability to change any nuclear property both quickly and easily.  In this work we have used the three built-in astrophysical modules in order to highlight the fact that with r-Java 2.0 for the first time anyone can run a comparison of proposed astrophysical r-process sites with their dynamical evolution and ejecta conditions.  R-Java 2.0 goes beyond existing codes by including the option to define a custom density evolution profile in order to study other possible r-process scenarios.  As many of the nuclear inputs used in r-Java 2.0 are only theoretically known, the results from rare-isotope beam facilities and upcoming sensitivity studies will provide much needed experimental insight into  r-process nuclei \\citep{dillman,hosmer,vanscelt}.  The flexibility of r-Java 2.0 allows for researchers to easily include experimental results in order to test their impact on the r-process.  The areas of development that we are undertaking for the next release of r-Java are;  the inclusion of neutrino-induced and spallation reactions, the development of a charged-particle reaction network module and adding the ability to study nuclear isomers.   \\onecolumn  \\clearpage \\begin{figure} \\includegraphics[scale = 0.75]{HEW.pdf} \\caption{\\textbf{High-entropy wind.}  See section \\ref{hewRes} for details. \\textit{Top:}  The black solid line denotes an entropy (S) of 140 and expansion velocity ($v_{\\rm exp}$) of 3$\\times10^4$ km s$^{-1}$.  The red dashed line denotes S = 175 and $v_{\\rm exp}$ = 1.5$\\times10^4$ km s$^{-1}$.  The green dotted line denotes S = 270 and $v_{\\rm exp}$ = 3.8$\\times10^3$ km s$^{-1}$.  NSE determined initial abundances.  \\textit{Middle:} Same as top panel with $\\alpha$-rich freeze-out initial abundances.  \\textit{Bottom:}  The black solid line denotes a neutron-to-seed ratio of 25, the red dashed line a neutron-to-seed ratio of 75 and the green dotted line a neutron-to-seed ratio of 100.  Each simulation started from $\\alpha$-rich freeze-out abundances.} \\label{HEWcomp} \\end{figure}      \\clearpage \\begin{figure} \\includegraphics{NSM_rhos.pdf} \\caption{Using the neutron star merger module the density evolution of ejecta with two different polytropic indices (n) are compared.  The black solid line denotes relativistically degenerate matter (n = 1) and the red dashed line represents non-relativistically degenerate matter (n = 3). } \\label{NSMrho} \\end{figure}  \\clearpage \\begin{figure} \\includegraphics[scale = 0.75]{NSM_compPhys.pdf} \\caption{\\textbf{Neutron star merger.}  Final r-process abundances are compared for varying physical parameters.  See section \\ref{NSMphysRes} for details of simulations.  \\textit{Top:} Varying polytropic indexes (n); n = 1 (black solid line), n = 2 (red dashed line) and n = 3 (green dotted line).  \\textit{Middle:}  Varying initial densities ($\\rho_0$); $\\rho_0 =  2\\times10^{11}$g cm$^{-3}$ denoted by the black solid line, $\\rho_0 =  3\\times10^{11}$g cm$^{-3}$ by the red dashed line and $\\rho_0 =  4\\times10^{11}$g cm$^{-3}$ by the green dotted line.  \\textit{Bottom:}  Varying expansion velocity ($v_{\\rm exp}$); $v_{\\rm exp} = 10^2$km s$^{-1}$ denoted by the black solid line, $v_{\\rm exp} = 10^3$km s$^{-1}$ by the red dashed line and $v_{\\rm exp} = 10^4$km s$^{-1}$ by the green dotted line   } \\label{NSMcompPhys} \\end{figure}       \\clearpage \\begin{figure} \\includegraphics[scale = 0.75]{NSM_CompInitial.pdf} \\caption{\\textbf{Neutron star merger.} The final nuclei abundances are plotted for different initial electron fractions ($Y_{\\rm e,0} = 0.1$ denoted by the black solid line, $Y_{\\rm e,0} = 0.15$ by the red dashed line and $Y_{\\rm e,0} = 0.2$ by the green dotted line.  See Sect. \\ref{NSMiniRes} for details on initial physical parameters.  \\textit{Top:}  Initially only neutron, protons and $\\alpha$-particles present.  \\textit{Middle:}  An initial seed of iron group isotopes.  \\textit{Bottom:} The initial abundances determined using NSE. } \\label{NSMcompInit} \\end{figure}      \\clearpage \\begin{figure} \\includegraphics[scale = 0.75]{initialNSE_compMejZeta.pdf} \\caption{\\textbf{Quark nova.} Comparison of the effect on r-process yield of varying the percentage of quark nova energy transformed into kinetic energy of the ejecta.  In each panel the black solid line denotes 1$\\%$, the red dashed line denotes 2$\\%$ and the green dotted line 5$\\%$.  For each panel the initial nuclei abundance distribution is determined by NSE.  \\textit{Top:}  Quark nova ejecta of $10^{-3}$M$_{\\odot}$.  \\textit{Middle:}  Quark nova ejecta of $10^{-4}$M$_{\\odot}$. \\textit{Bottom:}  Quark nova ejecta of $10^{-5}$M$_{\\odot}$. } \\label{QNcompMejZeta} \\end{figure}  \\clearpage \\begin{figure} \\includegraphics[scale = 0.75]{QN_compInitial.pdf} \\caption{\\textbf{Quark nova.}  Comparison of the final r-process abundance yield of different masses of quark nova ejecta ($10^{-3}$M$_{\\odot}$ denoted by the black solid line, $10^{-4}$M$_{\\odot}$ by the red dashed line and $10^{-5}$M$_{\\odot}$ by the green dotted line).  See section \\ref{QNiniRes} for details of simulations.  \\textit{Top:}  Initially only neutrons, protons and $\\alpha$-particles.  \\textit{Middle:} Initially a seed of iron group isotopes.  \\textit{Bottom:}  Initial abundance determined by NSE.} \\label{QNcompInitial} \\end{figure}     \\clearpage \\begin{figure} \\includegraphics[scale = 0.85]{CompSolar.pdf} \\caption{In each panel the superposition of the r-process yields (red solid line) from each astrophysical site is compared to the observed solar r-process abundances (black crosses).  The constituent components are as well plotted for each site.  \\textit{Top:} The r-process abundances from high-entropy wind events; $Y_{\\rm n} / Y_{\\rm seed} = 25$ (green dash-dot line), $Y_{\\rm n} / Y_{\\rm seed} = 50$ (blue dashed line) and  $Y_{\\rm n} / Y_{\\rm seed} = 100$ (black dotted line). \\textit{Middle:}  The r-process abundances from neutron star merger events; $Y_{\\rm e}  = 0.3$ and $\\rho = = 3 \\times 10 ^{11}$g cm $^{-3}$ (green dash-dot line), $Y_{\\rm e}  = 0.1$ and $\\rho = 8 \\times 10 ^{11}$g cm $^{-3}$ (blue dashed line).  \\textit{Bottom:}   The r-process abundances from quark novae; $m_{\\rm ej} = 10^{-4} M_{\\odot}$ and $\\zeta = 1\\%$ (green dash-dot line), $m_{\\rm ej} = 10^{-4} M_{\\odot}$ and $\\zeta = 5\\%$ (black dotted line) and $m_{\\rm ej} = 10^{-5} M_{\\odot}$ and $\\zeta = 1\\%$ (blue dashed line).  } \\label{sol} \\end{figure}    \\clearpage \\begin{figure} \\includegraphics{allSites.pdf} \\caption{All sites scaled by predicted rates compared to solar abundances. } \\label{allSites} \\end{figure}  \\clearpage \\begin{table} \\caption{Quark Nova Ejecta Parameters} % \\label{qnTable} % \\centering % \\begin{tabular}{c | c c c c} % \\hline\\hline % Mass & T$_{0}$ & $\\rho_0$ & Y$_{\\rm e,0}$ & $\\tau$ \\\\ %  [M$_{\\odot}$]& [10$^9$K] &  [10$^{11}$g cm$^{-3}$] & &[$\\mu$s] \\\\ \\hline % 10$^{-5}$ & 2.4 & 1.0 & 0.18 & 0.1 \\\\ % 10$^{-4}$ & 3.8 & 4.0 & 0.13 & 0.5 \\\\ 10$^{-3}$ & 7.5 & 30 & 0.07 & 6.0 \\\\ \\hline % \\end{tabular} \\end{table}  \\clearpage \\appendix"
    },
    "1402/1402.7016_arXiv.txt": {
        "abstract": "White-light observations by the Solar Dynamics Observatory's  Helioseismic and Magnetic Imager of a loop-prominence system occurring in the aftermath of an X-class flare on 2013 May 13 near the eastern solar limb show a linearly polarized component, reaching up to $\\sim$20\\% at an altitude of $\\sim$33 Mm, about the maximal amount expected if the emission were due solely to Thomson scattering of photospheric light by the coronal material. The mass associated with the polarized component was 8.2$\\times$10$^{14}$ g. At 15 Mm altitude, the brightest part of the loop was 3($\\pm$0.5)\\% linearly polarized, only about 20\\% of that expected from pure Thomson scattering, indicating the presence of an additional unpolarized component at wavelengths near Fe I (617.33 nm), probably thermal emission. We estimated the free electron density of the white-light loop system to possibly be as high as 1.8$\\times$10$^{12}$ cm$^{-3}$. ",
        "introduction": "\\citet{Martinez2014} (thereafter Paper I) have already reported on the observations by the Helioseismic and Magnetic Imager \\citep[HMI,][]{Schou2012b,Scherrer2012} of the Solar Dynamics Observatory \\citep[SDO,][]{Pesnell2012} of coronal emission from two flares occurring on 2013 May 13. Both of these also showed white-light (WL) footpoint sources at the level of the photosphere. The gradual coronal emissions can be identified as visual counterparts of the classical loop-prominence system, but were brighter than expected and possibly seen in the continuum rather than line emission, as inferred from high-resolution HMI spectra. In this interpretation, the coronal sources detected by HMI in these flares represent flare loops, initially heated to X-ray temperatures, and detected in the process of cooling. The authors found the HMI flux to exceed the radio/X-ray interpolation of the bremsstrahlung produced in the flare soft X-ray sources by at least one order of magnitude, implying the participation of cooler sources that could produce free-bound continua and possibly line emission detectable by HMI.  \\begin{figure*}[ht!] \\centering \\includegraphics[width=16.6cm]{Fig1.eps} \\caption{The gradual-phase sources, with image times: 16:25:22.7 (WL), 16:25:28.1 (160 nm), and 16:25:21.5 (19.3 nm). The AIA 19.3 nm image shows loop-shaped absorption features, corresponding loosely to the WL and UV loops. The purple contours are from RHESSI observations of thermal X-rays, and the red line is the 20\\% contour level of the HMI image. Reproduction of Fig. 5 from Paper I, with the addition of a green box used to accumulate the time profiles displayed in Figure~\\ref{fig:HMIlc}.} \\label{fig:FigX} \\end{figure*}  Historically, the loop-prominence phenomenology played a major role in establishing the standard model of large-scale magnetic reconnection as a mechanism for the formation of flare loops projected against the corona, initially through observations in chromospheric lines \\citep[e.g.,][]{Svestka1972}. Observations of polarization in broad-band continuum by coronagraphs typically do not extend low enough to study the arcade development; for example, the Mk4 K-coronameter at Mauna Loa only observes above about 1.12 R$_\\odot$, some 80 Mm above the photosphere. Nevertheless some direct broad-band intensity observations in the lower corona have been reported (Hiei et al., 1992; Leibacher et al., 2004). These did not include the polarization signatures that HMI provides.  In this paper, we will concentrate on one of the two events discussed in Paper I, SOL-2013-05-13T16:01 (X2.8), and we present evidence of a Thomson-scattered component in the HMI emission, which allows us to discriminate between emission mechanisms contributing to the observed HMI emission. We describe the interpretation of the polarization signatures in the Appendix.  \\begin{figure}[ht!] \\centering \\includegraphics[width=7.5cm]{Fig2.eps} \\caption{ 2013 May 13 GOES X-ray (black: 1--8 \\AA\\ ; orange: 0.5--4 \\AA\\ ) and HMI Stokes I, Q, U, V time profiles. The HMI lightcurves are generated by averaging over all six filters the pixels in the green box of Figure~\\ref{fig:FigX}. Each pixel within the box was temporally median-filtered (with window size 3). No other data alterations were performed. } \\label{fig:HMIlc} \\end{figure} ",
        "conclusions": "\\label{sect:ccl} We have reported what we believe to be the first detection of  linearly polarized scattered white-light of an evolving flare loop system in the vicinity of the Fe I line, and probably the first unambiguous mass estimate of such a system, owing to the linear dependence on density and temperature-independence of Thomson scattering. It is only due to SDO/HMI's remarkable dynamic range and polarimetric capabilities that such a faint signature of Thomson scattering could be identified. This has enabled us to dynamically 1) identify the fraction of WL emission in a flare loop at the limb that is due to Thomson scattering, and 2) estimate the free electron content and the mass of the scattering sources, as well as likely limits to their free electron densities. Such an approach brings a powerful new diagnostic tool to the study of limb flares, and possibly the flare/CME connection as well."
    },
    "1402/1402.2858_arXiv.txt": {
        "abstract": "{We present deep \\ion{H}{i} observations of three compact high-velocity clouds (CHVCs).} {The main goal is to study their diffuse warm gas and compact cold cores. We use both low- and high-resolution data obtained with the 100 m Effelsberg telescope and the Westerbork Synthesis Radio Telescope (WSRT). The combination is essential in order to study the morphological properties of the clouds since the single-dish telescope lacks a sufficient angular resolution while the interferometer misses a large portion of the diffuse gas.} {Here single-dish and interferometer data are combined in the image domain with a new combination pipeline. The combination makes it possible to examine interactions between the clouds and their surrounding environment in great detail.} {The apparent difference between single-dish and radio interferometer total flux densities shows that the CHVCs contain a considerable amount of diffuse gas with low brightness temperatures. A Gaussian decomposition indicates that the clouds consist predominantly of warm gas.} {} ",
        "introduction": "\\label{sec:introduction}  High-velocity clouds (HVCs) are neutral atomic hydrogen (\\ion{H}{i}) clouds with radial velocities incompatible with simple Galactic rotation models \\citep{1997ARA&A..35..217W}. After their initial discovery by \\citet{1963CRAS..257.1661M}, they have been observed all around the Milky Way. \\citet{1999ApJ...514..818B} suggested that HVCs are distributed throughout the Local Group, but this theory has been disproven several times by searching for HVCs around other galaxies: all these observations have either failed to find any HVCs (\\citet{2004ApJ...610L..17P, 2007ApJ...662..959P}) or located them very close to their host galaxies \\citep{2004ApJ...601L..39T}. High-velocity clouds are found either as compact and isolated objects or as part of large complexes \\citep{1997ARA&A..35..217W}.  Recent studies propose three main hypotheses for the origin of HVCs \\citep{2004ASSL..312..341B}. First, HVCs consist of primordial gas that is accreted onto galaxies. It can be primordial gas flows from the filaments or have its origin in gaseous dark matter haloes. Second, HVCs originate from the tidal and ram-pressure interaction of the Milky Way Galaxy with dwarf galaxies. Third, HVCs have been formed as the result of a galactic fountain, i.e., by gas flows driven by supernovae. The most serious problem is the limited information on the HVC distances \\citep{2000ASPC..218..407V, 2004ASSL..312..195V, 2005A&A...440..775K, 2008ASPC..393..179B}.  Investigations of the physical conditions of HVCs indicate that some show signs of interaction with the Galactic halo gas \\citep{2000A&A...357..120B, 2001A&A...370L..26B, 2005A&A...432..937W, 2011MNRAS.418.1575P, 2012A&A...547A..12V}.  The aim of our study is to investigate the morphological properties of compact high-velocity clouds (CHVCs) as well as their radial velocity and line width. In the past, the existence of two different gradients either in density \\citep{2011MNRAS.418.1575P} or velocity \\citep{2000A&A...357..120B} were required within the cloud for the designation head-tail (HT). The objects studied in the present paper appear to conform to both these criteria. The appearance of the HT structure is interpreted as follows. As the cloud moves through the hot and ionized halo, part of the cloud is compressed. This higher density area constitutes the head of the cloud. Part of the gas is stripped off the cloud to form a less dense and thin tail structure that follows the head with velocities lower than those of the cloud bulk motion \\citep{2002A&A...391..713K, 2009ApJ...698.1485H}.  The cloud HVC\\,125+41--207 is prototypical for interacting HVCs denoted accordingly as head-tail (HT) HVCs. \\citet{2000A&A...354..853B} used the Westerbork Synthesis Radio Telescope (WSRT) to study the small--scale structure of compact HVCs (CHVCs). Towards HVC\\,125+41--207, they discovered an extremely narrow \\hi line unresolved by their spectral resolution of $2.47\\,{\\rm km\\,s^{-1}}$. The peak brightness temperature of this line is exceptionally high with $T_{\\rm B} = 75$\\,K indicating the existence of a cold and dense cold neutral medium (CNM) core. Their maps reveal a highly structured \\hi CNM distribution. \\citet{2000A&A...357..120B} used Effelsberg observations of HVC\\,125+41--207 to deduce that this dense core is embedded within a warm neutral medium (WNM) envelope. Moreover, they found evidence that the WNM shows signs of deceleration not only towards the cloud's tail but also along the rims. Proportional to this deceleration of the radial velocity component, the WNM gets warmer (see their Fig.\\,2). Only the combined analysis of the WSRT and Effelsberg data allowed the deduction of such a homogenous view of HVC\\,125+41--207 as an interacting CHVC.  To investigate the gaseous structures of HVCs and their dynamics accurately, the combination of radio interferometric and single-dish data is essential. Radio interferometers are insensitive to structures on the scale of tens of arc minutes and beyond, the WNM. Single-dish observations are unable to resolve the small-scale structure of the CNM. The combination of \\ion{H}{i} single-dish and interferometric observations provides the possibility of studying HVCs in the necessary detail.  In this paper, we present high-resolution \\ion{H}{i} observations of three CHVCs, using the 100 m Effelsberg telescope and the WSRT. Section\\,\\ref{sec:observations} presents the details of the data, and the \\ion{H}{i} observations are introduced. Section\\,\\ref{subsec:method} deals with the method that has been used to combine these two data sets. Section\\,\\ref{sec:CHVCs} discusses the physical and morphological properties of the clouds and the results of the Gaussian decomposition of the integrated spectral lines. Section\\,\\ref{sec:conclusions} summarizes our findings and gives an outlook on future work. ",
        "conclusions": "\\label{sec:conclusions}  We present deep integrated observations of the warm neutral medium for three CHVCs using the 100 m Effelsberg telescope and high-resolution observations of the more compact regions from the WSRT, as well as the results of their combination. Our approach shows that the combination in the image domain meets the expectations. The critical step in the algorithm is the regridding, where a flux inconsistency can occur as a result of interpolation inaccuracies.  The combination results demonstrate the importance of the zero-spacing correction regarding determination of the physical and morphological properties of the objects. Here in particular the warm neutral medium is of major interest. High-velocity clouds in higher ambient pressure environments might show complex spatial and spectral structures.  The results of the Gaussian decomposition suggest strongly that the three clouds mostly consist of WNM. The analysis of \\citet{2011A&A...533A.105W} also confirms the lack of cold components in HVCs in the complex known as Galactic centre negative (GCN). However, high Doppler temperatures above $10^4$ K are an indication of the existence of turbulence within the clouds. The results demonstrate that the gas gets warmer or more turbulent in the tail region and at the edges of the clouds. This also suggests that the warm gas is floating away to the direction of the tail, which is an indication of existing ram pressure. Another remarkable aspect is the decreasing velocity gradient at the regions with lower column densities. The relatively slow movement of the clouds also reveals the lack of cold cores in the clouds, and the measured high Doppler temperatures demonstrate that the clouds are not in equilibrium.  The linear approach used in the image domain is easy to implement and  is not very CPU-intensive. It could be very helpful for computation of huge amounts of data coming from telescopes like ASKAP and WSRT/APERTIF."
    },
    "1402/1402.5363_arXiv.txt": {
        "abstract": "We have conducted a survey of 17 wide ($>$ 100 AU) young binary systems in Taurus with the Atacama Large Millimeter Array (ALMA) at two wavelengths. The observations were designed to measure the masses of circumstellar disks in these systems as an aid to understanding the role of multiplicity in star and planet formation. The ALMA observations had sufficient resolution to localize emission within the binary system. Disk emission was detected around all primaries and ten secondaries, with disk masses as low as 10$^{-4}$ M$_{\\odot}$.  We compare the properties of our sample to the population of known disks in Taurus and find that the disks from this binary sample match the scaling between stellar mass and millimeter flux of F$_{mm} \\propto$ M$_{\\ast}^{1.5-2.0}$ to within the scatter found in previous studies. We also compare the properties of the primaries to those of the secondaries and find that the secondary/primary stellar and disk mass \\emph{ratios} are not correlated; in three systems, the circumsecondary disk is more massive than the circumprimary disk, counter to some theoretical predictions. ",
        "introduction": "\\label{intro}  Most stars are formed in binary or multiple systems and remain in such systems for their main sequence lifetimes \\citep [e.g.][]{mon07}. Therefore, understanding the causes and effects of multiplicity is an essential ingredient of complete models of both star and planet formation. Circumstellar disks play a crucial role in both processes, tracing effects of different binary formation mechanisms, providing conduits for material to accrete onto the stars, and serving as the reservoir of raw material for planet formation.  At a given point in time, the distribution of observed disk masses is a function of the initial disk masses and disk evolution.  For multiple systems, dynamical interactions between the stars, the circumstellar disks, and any circumbinary material will also impact both the disk formation and evolution. Models of binary star formation by \\citet{bat00} predict that the circumprimary disk, i.e. the disk around the more massive star, will have more mass than the circumsecondary disk; however, these models do not follow the viscous evolution of the disk after the formation stage. Observations to date largely support the prediction of a more massive circumprimary disk,  although the sample of systems observed is relatively small and generally comprise only the brightest sources. \\citet{jen03} found that the primary star had the most massive disk in all four young binaries they observed; indeed in only one system was the secondary's disk detected at all, despite most of the secondaries showing signs of accretion. More recent work by \\citet{har12} has expanded the number of observed binary systems and also found that when both components were detected, the primary had higher flux, but with sensitivity levels of a few mJy, many secondaries remained undetected.  Planet formation in these systems may also be impacted as models of the interactions of binary stars with their associated disks predict that the disks will be truncated somewhere between 0.2 and 0.5 times the binary separation, depending on the eccentricity of the system \\citep{art96}.   However, these models do not address the surface density and evolution of the remaining disk material. If the secondary disks retain roughly the same surface density as the inner parts of disks around single stars, then they may still retain enough mass to form planets.  Previous observations have not had the sensitivity to distinguish between disks that are simply truncated, and those that have been significantly depleted by further accretion. Disk models show that truncation effects can affect the observed flux for separations up to a few hundred AU \\citep{jen96}.  The essential question for planet formation, then, is whether or not the disks around individual components of close binary stars are similar to the inner regions of disks around single stars. Early observations demonstrated that the unresolved millimeter emission, which traces the dust in the outer regions of the disk, is indeed reduced, consistent with truncation \\citep{bec90,ost95,jen96}. But most observations of binaries with separations in the ranges of 50-100 AU have yielded upper limits rather than detections, and indeed only about half of all young binaries in Taurus have been detected at all at millimeter wavelengths, despite the fact that many more than half of them were detected by IRAS at 60 $\\mu$m. With the advent of ALMA observations, which provide a substantial increase in sensitivity at the required resolution, it is now possible to reach much lower disk surface densities, and possibly to detect very low mass protoplanetary disks.  To address these issues, we have obtained ALMA Cycle 0 observations of 17 young binary systems in Taurus for which the components can be resolved.  In \\S \\ref{sample} we describe the sample selection and properties, in \\S \\ref{obs} we describe the ALMA observations and data reduction, in \\S \\ref{results} we present the results,  and we give our conclusions in \\S \\ref{conclusions}.   \\begin{deluxetable*}{llllll}[b!] \\tabletypesize{\\footnotesize} \\tablewidth{0pt} \\tablecaption{Taurus Binary Sample} \\tablehead{ \\colhead{Source Name}           & \\colhead{Additional names}      & \\colhead{Separation}  & \\colhead{mm flux\\tablenotemark{a}} & \\colhead{Spectral types\\tablenotemark{b}} & \\colhead{Previous mm} \\\\ & &  \\colhead{(AU)}  & \\colhead{(mJy)} & & \\colhead{detection\\tablenotemark{c}} } \\startdata FV Tau & & 101 & 48$\\pm$5 & K5/K6 & 1,2\\\\ HBC 387 & FV Tau/c & 104 & $<$25 & M2.5/M3.5 & --\\\\ FQ Tau & & 110 & 28$\\pm$7 & M3/M3.5 & 1\\\\ UY Aur & & 120 & 102$\\pm$6 & M0/M2.5 & 1,2\\\\ FX Tau & & 130 & 17$\\pm$3 & M1/M4 & 1\\\\ HBC 411 & CoKu Tau/3 & 290 & $<$8 & M1/M4.5 & --\\\\ IRAS 05022+2527 & CIDA 9 & 320 & 71$\\pm$7 & K8/M1.5 & 1\\\\ HK Tau & & 340 & 130$\\pm$2 & M0.5/M2 & 1,2\\\\ IT Tau & & 340 & 22$\\pm$3 & K3/M4 & 1,2\\\\ DK Tau & & 350 & 80$\\pm$10 & K8/M1 & 1\\\\ GK Tau & & 340 & 33$\\pm$7 & K7 & 1\\\\ HN Tau & & 430 & 29$\\pm$3 & K5/M4 & 1 \\\\ V710 Tau & & 450 & 152$\\pm$6 & M0.5 & 1\\\\ IRAS 04113+2758 & MHO 1/2& 550 & 380$\\pm$3  & M2.5/M2.5 & 1,2\\\\ IRAS 04298+2246 & JH 112 & 920 & 30$\\pm$10 & K6/M8.5 & 1\\\\ HO Tau & & 970 & 44$\\pm$6 & M0.5 & 1 \\\\ DS Tau & & 990 & 39$\\pm$4 & K5 & 1\\\\ \\enddata \\label{tab:sample} \\tablenotetext{a}{Millimeter fluxes are single-dish fluxes at 850 $\\mu$m taken from \\citet{and05} except for HK Tau \\citep[850 $\\mu$m, interferometry]{har12} and IRAS 04113+2758 \\citep[1.3 mm, interferometry]{har12}  } \\tablenotetext{b}{Spectral types from \\citet{and13}} \\tablenotetext{c}{1 = primary, 2 = secondary} \\end{deluxetable*}  \\section {Sample} \\label{sample}  We selected targets from a single star formation region, Taurus (distance $\\sim$ 140 pc), so that effects such as age and cluster environment are kept constant as much as possible.   Taurus is ideal in having a significant population of YSOs that have evolved into the disk-only state (with no remaining envelope) and is very well studied, containing both a well known set of single stars with disks for comparison and a significant population of binaries where both stellar components have been characterized in the optical or near-infrared. We started with the list of known Taurus binaries \\citep{and05,kra09} and selected those with separations in the range of 0\\farcs7 (100 AU) to 10\\arcsec\\ (1400 AU). The inner cutoff was selected such that the two components could be clearly resolved with the resolution offered in Cycle 0, while the outer cutoff was chosen to ensure that the systems are likely to be physically associated.    We eliminated systems classified as Class I from their spectral energy distributions, as these systems often contain substantial envelope emission that must be disentangled from the disk emission, and we eliminated systems with no active accretion signatures that had not been previously detected at millimeter wavelengths.  The resulting sample contains 17 systems (Table \\ref{tab:sample}) and includes all Class II Taurus binaries with separations of 100 to 1400 AU from \\citet{and05} and \\citet{kra09}. Higher-order multiple systems were excluded where known at the time of our sample selection, although two were observed (see notes in \\S \\ref{individ}). ",
        "conclusions": "  \\begin{itemize} \\item The majority of our new detections were for secondary sources and for these wide binaries ($> 100$ AU), the secondary disk fraction is somewhat higher than shown in previous studies.  We found 11 of the 15 bona fide young stellar binaries have disks with masses  $\\geq 10^{-4}$M$_{\\odot}$ around both stars, while \\citet{jen03} detected a circumsecondary disk in 1 of 4 systems and \\citet{har12} detected a circumsecondary disk in 6 of 12 systems where the components were resolved. There is significant overlap in the samples between these studies and the new detections are primarily due to the higher sensitivity of the ALMA observations. The newly detected primary disk masses and most of the secondary disk masses are considerably smaller than the minimum mass solar nebula, but this is not surprising given that the host stars are generally less than 1 M$_{\\odot}$.  While it may be difficult for massive planets to form in these less massive disks, models of core accretion around lower-mass stars show that they may be able to form cores for lower-mass planets \\citep[e.g.][]{lau04} and if the disk mass has evolved, larger planets may have formed earlier when the disk was more massive.  \\item In two binary systems the secondary disk has a higher mm flux than the primary disk.  This has not been seen in previous, smaller surveys and is counter to predictions of formation models where the infalling material is directly accreted onto the primary or secondary disk as opposed to accreting onto a circumbinary structure. This result could be explained by faster dissipation of the primary disk, which has been shown to be a function of stellar mass.  \\item For this sample of wide binaries,  the secondary/primary disk mass ratio is not correlated with the secondary/primary stellar mass ratio.  This suggests that for these binary systems, any environmental factors shared between the two components that could affect the initial disk mass and disk evolution are not the dominant factor in determining the range of disk masses for a given stellar mass.  \\end{itemize}  From these conclusions, it is clear that binaries do not follow a simple pattern of primary/secondary disk mass distribution; therefore, care should be taken when assigning flux to components in unresolved systems."
    },
    "1402/1402.4479_arXiv.txt": {
        "abstract": "{} {We investigate whether the cylindrical (galactocentric) radial velocity gradient of $\\sim-3\\kmseckpc$, directed radially from the Galactic center and recently observed in the stars of the solar neighborhood with the \\emph{RAVE} survey, can be explained by the resonant effects of the bar near the solar neighborhood.} {We compared the results of test particle simulations of the Milky Way with a potential that includes a rotating bar with observations from the \\emph{RAVE} survey. To this end we applied the \\emph{RAVE} selection function to the simulations and convolved these with the characteristic \\emph{RAVE} errors. We explored different ``solar neighborhoods\" in the simulations, as well as different bar models.} {We find that the bar induces a negative radial velocity gradient at every height from the Galactic plane, outside the outer Lindblad resonance and for angles from the long axis of the bar compatible with the current estimates. The selection function and errors do not wash away the gradient, but often make it steeper, especially near the Galactic plane, because this is where the \\emph{RAVE} survey is less radially extended. No gradient in the vertical velocity is present in our simulations, from which we may conclude that this cannot be induced by the bar.} {} ",
        "introduction": "Many of the past efforts in modeling the mass distribution of the Milky Way have assumed that the Galaxy is axisymmetric and in a steady state. However, there is a wealth of evidence that these assumptions are not really valid. The two most important deviations from axisymmetry are the spiral arms and the bar. These features are not only apparent as non-axisymmetric density enhancements, but they also have long-range gravitational effects. In particular, the bar modifies the kinematics of the outer parts of the Galactic disks, far beyond its extension, through resonant interactions.  That the velocity distribution of stars very near to the Sun is not smooth (as one would expect in a steady state axisymmetric system), but instead rich in substructures, has been established observationally thanks to data from the \\emph{Hipparcos} satellite and other surveys (\\citealt{Dehnen1998,Famaey2005,Antoja2008}). Several authors have explained these substructures as being due to orbital resonant effects of the bar (\\citealt{Dehnen2000, Fux2001}), of the spiral arms (\\citealt{Mayor1970,DeSimone2004}; \\citealt{Antoja2011}), or both (\\citealt{Antoja2009,Quillen2011}).  Using data from the \\RAVE survey (\\citealt{Steinmetz2006}), \\cite{Antoja2012} discovered that some of the kinematic substructures detected in the vicinity of the Sun can be traced further, both on and above the plane of the Galaxy, up to $\\sim 0.7\\Kpc$.  But \\emph{RAVE} also made it possible to discover large scale streaming motions. \\cite{Siebert2011grad} (in the rest of the paper S11) used a sample of $213,713$ stars to discover a gradient in the mean galactocentric radial velocity that decreases outward with Galactic radius. S11 show that this gradient was also present when using only the $29,623$ red clump stars in their sample, whose distances are more accurate. \\cite{Siebert2012} modeled the gradient as caused by a long-lived spiral pattern. \\cite{Williams2013} (hereafter W13) studied the 3D velocity distribution of red clump stars in \\emph{RAVE} in detail, confirmed the existence of the radial velocity gradient and also discovered a more complicated vertical velocity distribution than expected, attributing it to secular phenomena in the Galaxy. \\cite{Faure2014} generalized to 3D the model for the spiral arms presented in \\cite{Siebert2012}, which now also depends on the distance from the Galactic plane. This model nicely predicts a behavior for the mean vertical velocity that is similar to what is observed in W13 (i.e., resembling ``rarefaction-compression'' waves), together with the radial velocity gradient.  On the other hand, \\cite{Monari2013} (hereafter M13), used 3D test particle simulations to show that the gravitational effects of the bar can significantly affect the kinematics of stars near the Sun, even at distances from the Galactic plane up to at least $z\\sim1\\Kpc$ for the thin disk and $z\\sim2\\Kpc$ for the thick disk. These results imply that some of the substructures found in \\cite{Antoja2012} could also be caused by the bar.  In this paper we investigate an alternative explanation for the observed radial velocity gradient, beyond that caused by the spiral arms, by suggesting that it can be created by the bar. To do so, we compare the results of the test particle simulations in M13 with the \\emph{RAVE} data.  The paper is organized as follows. In Sect.~\\ref{sect:sim} we summarize the salient characteristics of the simulations from M13, and in Sect.~\\ref{sect:sf} we describe how we apply the \\RAVE selection function and error convolution to them, to mimic a \\RAVE catalog. In Sect.~\\ref{sect:res} we present the results. In Sect.~\\ref{sect:expl} we explain how the bar can create a radial velocity gradient as observed. In Sect.~\\ref{sect:disc} we discuss the similarities between our results and the ones in W13 and in Sect.~\\ref{sect:concl} we conclude. ",
        "conclusions": "\\label{sect:concl} In this work we have proposed a new explanation for the recent discovery (S11, W13) of a negative $R$ gradient of the (galactocentric) radial velocity. We found that the bar can create a negative gradient if the Sun is placed just outside the outer Lindblad resonance and at angles from the long axis of the bar similar to the current estimates from the literature. The velocity gradients become steeper when increasing the angle from the bar and also for the Long Bar model. On the other hand, in the less massive bar case they become shallower. Moreover, such gradients do not depend strongly on the height from the Galactic plane. This happens because the bar affects the kinematics of the Galaxy almost in the same way from $z=0$ to $z\\sim2\\Kpc$, as explained in M13. Because of this, the bar provides a natural mechanism for the observed gradients at different heights.  We compared the 3D test particle simulations presented in M13 with the findings of \\emph{RAVE}, after applying the \\emph{RAVE} selection function and proper error convolution. The gradients exist in our simulations for all bar parameters and positions of the Sun explored (all outside the outer Lindblad resonance). These gradients are never completely washed out by the selection function and the errors, but rather they are enhanced in some cases. In fact, the gradients in the solar neighborhood spheres considered are in general shallower than those observed in the Milky Way, but the selection function can enhance them to the level of $\\sim3-4\\kmseckpc$ (as e.g., happens for all the studied simulation slices with $|z|<0.5\\Kpc$ and the Sun centered at $R=8\\Kpc$).  However, none of the models that we explored in this work accurately describes the behavior in \\RAVE of $\\avvR$ at every $z$: the gradients are too shallow for $z<0$. Some models resemble \\RAVE for $z>0$, especially our default bar case at $R=8\\Kpc$ and $\\phi\\leq-20\\degr$. We conclude from this that the bar should at least contribute to the negative gradient observed, for position angles with respect to the bar $\\phi < 0$ and for locations of the Sun near but outside the outer Lindblad resonance ($R > \\ROLR$).  Furthermore, our simulations do not show any kind of vertical velocity gradient as seen in the data for \\RAVE by W13. This result is consistent with the distribution function of the simulated disks being an even function of $v_z$. On the other hand, the recent paper by \\cite{Faure2014} shows that a 3D model for spiral arms is successful in reproducing radial and vertical velocity gradients similar to those observed in W13. In reality both effects of bar and spiral arms probably coexist and shape the velocity distribution of the solar neighborhood. However, while in the case of the bar the slope of the radial velocity gradient depends significantly on the angular location of the observer in the Galaxy, in the case of tightly wound spirals the angle is much less important (Fig.~6 and 7 in \\citealt{Faure2014}). Future observations of the Galactic disk (e.g., obtained with the \\emph{Gaia} satellite) are expected to be sufficiently extended to distinguish whether the main cause of the radial velocity gradient is the bar or the spiral arms.  A natural future development of this work is to fit the kinematics of the extended solar neighborhood with the analytic predictions from the bar perturbation theory, in the same fashion as in \\cite{Siebert2012} for the spiral arms, in order to retrieve the best fit values for the bar pattern speed, bar angle, and bar strength."
    },
    "1402/1402.6336_arXiv.txt": {
        "abstract": "Gamma Ray Bursts (GRBs) are characterized by ultra-relativistic outflows, while supernovae are generally characterized by non-relativistic ejecta. GRB afterglows decelerate rapidly usually within days,  because their low-mass ejecta rapidly sweep up a comparatively larger mass of circumstellar material. However supernovae, with heavy ejecta, can be in nearly free expansion for centuries. Supernovae were thought to have non-relativistic outflows except for few relativistic ones accompanied by GRBs. This clear division was blurred by SN 2009bb, the first supernova with a relativistic outflow without an observed GRB. Yet the ejecta from SN 2009bb was baryon loaded, and in nearly-free expansion for a year, unlike GRBs. We report the first supernova discovered without a GRB, but with rapidly decelerating mildly relativistic ejecta, SN 2012ap. We discovered a bright and rapidly evolving radio counterpart driven by the circumstellar interaction of the relativistic ejecta. However, we did not find any coincident GRB with an isotropic fluence of more than a sixth of the fluence from GRB 980425. This shows for the first time that central engines in type Ic supernovae, even without an observed GRB, can produce both relativistic and rapidly decelerating outflows like GRBs. ",
        "introduction": "\\subsection{Ordinary Supernovae} The optical lightcurves of supernovae have been well described by \\citet{1982ApJ...253..785A} as a nearly black body photosphere which recedes into the ejecta, heated by $\\gamma$-rays from nuclear decay and cooled by rapid expansion with a characteristic velocity of $v\\sim 10^4$ km s$^{-1}$. The interaction of the ejecta with the circumstellar medium set up by the stellar wind of the progenitor star has been described by \\citet{1982ApJ...258..790C} using self similar solutions. Such solutions have also successfully described the radio emission from type Ic supernovae \\citep{1998ApJ...499..810C}. The combination of nearly free expansion leading to decreasing optical thickness from free-free or synchrotron self absorption (SSA) and decreasing magnetic fields produces a decreasing peak frequency (the frequency at which the peak in the radio spectrum occurs), but a nearly constant flux density at the peak frequency \\citep{1998ApJ...499..810C}. In the self-similar solution, this interaction produces a shockfront which expands in a powerlaw fashion, with $R \\propto t^m$, where $R$ is the radius of the shockfront, $t$ is the time since shock breakout and $m$ is known as the expansion parameter (also sometimes called the deceleration parameter). In the case of Type Ib/c supernovae, which generally have a relatively tenuous circumstellar medium, this interaction produces little deceleration, and $m$ is close to 1. The ejecta would slow down significantly only after encountering a mass of external medium comparable to the $\\sim1$ M$_\\odot$ of ejecta expected for Type I b/c supernovae. This Sedov time \\citep{1950RSPSA.201..159T} is expected to be $\\gtrsim 10^2$ years for supernovae. Therefore young supernovae are usually found in a Newtonian phase of nearly free expansion.  \\subsection{Gamma Ray Bursts} GRBs were discovered by  \\cite{1973ApJ...182L..85K} using the Vela satellites, designed for the detection of nuclear tests in space. Ultra-relativistic blast waves had already been described by fluid dynamical \\citet{1976PhFl...19.1130B} solutions before they were implicated in the production of GRBs \\citep{1986ApJ...308L..47G,1986ApJ...308L..43P}. The production of GRBs require the ejecta to have initial bulk Lorentz factors of $\\Gamma \\gtrsim 10^2$ \\citep{1999PhR...314..575P} in order to overcome the pair production opacity \\citep{1975NYASA.262..164R,1978Natur.271..525S}. Furthermore, even a small number of baryons in the initial ejecta can soak up most of the explosion energy available for $\\gamma$-rays \\citep{1990ApJ...365L..55S}. This is called baryon poisoning. So GRB jets must have $\\lesssim 10^{-6}$ M$_\\odot$ of relativistic ejecta \\citep{1999PhR...314..575P} after breakout from the stellar progenitor surface, so as not to be baryon poisoned. In a short time, this relatively small mass of ejecta encounters a larger mass of external matter. Therefore GRB afterglows are found in a relativistic but rapidly decelerating phase and have $\\Gamma \\lesssim 20$. The broad-lined (high velocity) type Ic SN 1998bw was discovered in the direction of GRB980425 \\citep{1998Natur.395..670G}, and was characterized by the bright radio emission from its relativistic ejecta which had $\\Gamma \\sim 2-3$ \\citep{1998Natur.395..663K}. This association between a Type Ic supernova and a long GRB was followed by other supernovae like that of GRB030329 with SN 2003dh \\citep{2003Natur.423..847H} and XRF060218 with SN 2006aj \\citep{2006Natur.442.1011P,2006Natur.442.1014S} also associated with GRBs. Some of these supernovae were marked by broad lines or by asphericity. Yet the definitive property, of this small subset of Ic supernovae, that let them produce relativistic outflows remains elusive.  \\subsection{Search for relativistic ejecta} This evidence, for a supernova-GRB connection, inspired a systematic search for relativistic ejecta from nearby type Ic supernovae using radio observations, leading to the discovery of SN 2009bb by \\citet{2010Natur.463..513S} with $\\gtrsim10^{49}$ erg of energy in radio emitting relativistic ejecta. Like GRBs, SN 2009bb had relativistic ejecta \\citep{2010Natur.463..513S} but this ejecta continued to be in nearly free expansion for $\\sim 1$ year \\citep{2011NatCo...2E.175C,2010ApJ...725....4B}, leading to the suggestion, that unlike GRBs, it is baryon loaded \\citep{2011ApJ...729...57C}. It has been suggested by \\citet{2011NatCo...2E.175C} that such engine-driven relativistic supernovae can even accelerate ultra-high-energy cosmic rays. Such a baryon loaded fireball has also been implied in PTF11agg \\citep{2013ApJ...769..130C}. A rapid decline in flux density, faster than $\\propto t^{-1}$ at a given radio frequency, or a decreasing peak flux density may signal deceleration of the ejecta. Recently the rapidly declining PTF12gzk was observed with $\\lesssim10^{46}$ erg of energy in fast ejecta \\citep{2013ApJ...778...63H}. So far, however, highly energetic ($\\gtrsim10^{49}$ erg) relativistic ejecta combined with rapid deceleration has never been observed until now in a supernova unassociated with a GRB. ",
        "conclusions": "Even though no GRB counterpart was found for SN 2012ap, its radio afterglow shares remarkable characteristics with those of GRB associated supernovae; it has a relativistic ejecta component with relatively few baryons. While SN 2009bb had a relativistic outflow, it was clearly baryon loaded which is not the case here. A rapid radio decline has been observed in PTF12gzk which may also have had fast ejecta. But SN 2012ap had orders of magnitude more energy in relativistic ejecta, putting its energetics firmly in the class of GRB-associated supernovae. SN 2012ap pushes the boundaries of explosion parameters observed from stripped core progenitors with central engines (see Figures \\ref{vm_schem} and \\ref{momemass}). The parameters that separate the outflows, from GRB associated supernovae and ordinary Type Ib/c supernovae, according to \\citet{2011ApJ...729...57C} are the velocity and baryon loading of their fastest ejecta. According to \\citet{2014arXiv1402.6344M} the duration of central engine activity drives the diversity of explosion outcomes. By bridging the gap between ordinary supernovae and GRB associated supernovae in terms of its high velocity and low ejecta mass, SN 2012ap demonstrates the role of CEDEXs in understanding the supernova-GRB connection."
    },
    "1402/1402.1499_arXiv.txt": {
        "abstract": "We study the relationship between the stability level of late-type galaxy disks and their star-formation activity using integral-field gaseous and stellar kinematic data.  Specifically, we compare the two-component (gas$+$stars) stability parameter from \\citeauthor{2011MNRAS.416.1191R} ($\\qrw$), incorporating stellar kinematic data for the first time, and the star-formation rate estimated from 21cm continuum emission.  We determine the stability level of each disk probabilistically using a Bayesian analysis of our data and a simple dynamical model.  Our method incorporates the shape of the stellar velocity ellipsoid (SVE) and yields robust SVE measurements for over 90\\% of our sample.  Averaging over this subsample, we find a meridional shape of $\\sigma_z/\\sigma_R = 0.51^{+0.36}_{-0.25}$ for the SVE and, at 1.5 disk scale lengths, a stability parameter of $\\qrw = 2.0\\pm 0.9$.  We also find that the disk-averaged star-formation-rate surface density ($\\Ssfre$) is correlated with the disk-averaged gas and stellar mass surface densities ($\\Sge$ and $\\Sse$) and anti-correlated with $\\qrw$.  We show that an anti-correlation between $\\Ssfre$ and $\\qrw$ can be predicted using empirical scaling relations, such that this outcome is consistent with well-established statistical properties of star-forming galaxies.  Interestingly, $\\Ssfre$ is not correlated with the gas-only or star-only \\citeauthor{1964ApJ...139.1217T} parameters, demonstrating the merit of calculating a multi-component stability parameter when comparing to star-formation activity.  Finally, our results are consistent with the \\citeauthor{2010ApJ...721..975O} model of self-regulated star-formation, which predicts $\\Ssfre/\\Sge\\propto\\Sse^{1/2}$.  Based on this and other theoretical expectations, we discuss the possibility of a physical link between disk stability level and star-formation rate in light of our empirical results. ",
        "introduction": "\\label{sec:intro}  Stars are formed by the collapse of gas.  In galaxy disks, a gas cloud should be gravitationally unstable if (1) it cannot adjust its internal pressure to balance the local gravitational pressure on timescales shorter than a free-fall time and (2) it occupies an area smaller than the scale on which differential rotation will shear it apart.  \\citet{1964ApJ...139.1217T} codified these concepts into a criterion for the stability of an infinitely thin, rotating, self-gravitating, fluid disk: \\begin{equation} Q = \\frac{\\kappa \\sigma}{\\pi G \\Sigma} > 1, \\label{eq:q} \\end{equation} where \\begin{equation} \\kappa^2 = 2\\frac{v_c}{R}\\left( \\frac{v_c}{R} + \\frac{\\partial v_c}{\\partial R}\\right) \\label{eq:efreq} \\end{equation} is the epicyclic frequency, $v_c$ is the circular speed of the potential, $\\sigma$ is the radial velocity dispersion, $\\Sigma$ is the mass surface density, and $G$ is the gravitational constant.  However, self-gravity is only one major player in the star-formation process, with chemodynamical processes, turbulence, and magnetic fields also significantly affecting the dynamics \\citep[see][and references therein]{2007ARA&A..45..565M}.  The relative importance of these physical properties to the star-formation law is a matter of ongoing debate.  Empirical studies of the star-formation law in disk galaxies have predominantly focused on observations of their gaseous components.  In a seminal article, \\citet[][hereafter \\citetalias{1998ApJ...498..541K}]{1998ApJ...498..541K} demonstrated that the star-formation rate per unit area ($\\Sigma_{\\rm SFR}$), or equivalently the time derivative of the stellar mass surface density ($\\Ssfr \\equiv \\Sigma_{\\rm SFR}$), is well correlated with the surface density of the hydrogen gas, $\\Sigma_{\\rm H} = \\sdhi+\\sdmh$, following the star-formation law suggested by \\citet{1959ApJ...129..243S}.  Considering the proportionality from equation \\ref{eq:q}, one might expect such a relation if star-formation is driven by self-gravity.  Owing much to the flood of relevant data, the quantitative details of the star-formation law in galaxy disks and its relation to the Kennicutt-Schmidt (KS) law \\citepalias[$\\Ssfr \\propto \\Sigma_{\\rm H}^{1.4}$;][]{1998ApJ...498..541K} have been greatly scrutinized.  This scrutiny has lead to a number of alternatives to this paradigm; see compilations by, e.g., \\citet{2008AJ....136.2782L} and \\citet{2013MNRAS.434.3389Z}.  Many of these alternatives involve consideration of gas properties, such as accounting for the dust-to-gas ratio \\citep{2013AJ....146...19L}, or use of specific gas tracers.  \\citet{2002ApJ...569..157W} and \\citet{2008AJ....136.2846B} have shown that $\\Ssfr$ is more tightly correlated with $\\sdmh$ than with $\\Sigma_{\\rm H}$.  Therefore, the star-formation law will also be affected by the ability of the interstellar medium to convert \\hone\\ to H$_2$ \\citep{2004ApJ...612L..29B, 2006ApJ...650..933B}, which can be related to the hydrostatic pressure in the disk plane \\citep{1989ApJ...338..178E, 1993ApJ...411..170E}.  As an inherently dynamical process, the star-formation law should also be influenced by relevant dynamical timescales, such as the local free-fall time \\citep{2012ApJ...745...69K}. \\citetalias{1998ApJ...498..541K} considered a star-formation law that incorporated the orbital timescale \\citep{1997ApJ...481..703S, 1997RMxAC...6..165E}, which is relevant in a scenario where dynamical processes in the disk (such as bars and spiral arms) are a primary driver of star formation.  Here, we explore the role of the disk stability level in the star-formation process.  Although the star-formation law may be written to show an explicit dependence on the {\\it gas} stability parameter \\citep[$Q_g$; see, e.g.,][]{2005ApJ...630..250K}, the most relevant assessment of the stability level includes both the gas and the stars \\citep{2005ApJ...626..823L}.  Indeed, \\citet{2003MNRAS.346.1215B} have shown that the two-component (gas$+$stars) stability parameter of 16 galaxies is a better estimator of the star-formation {\\em threshold} than one incorporating the gas alone.  In general, however, the effect of the stellar component on the star-formation {\\em law} is not well-understood:  \\citet{2003MNRAS.346.1215B} have also shown that a star-formation law that considers only the gas component (the KS law) is statistically indistinguishable from one proposed by \\citet{1994ApJ...430..163D} that incorporates the {\\em total} disk mass surface density.  More recently, \\citet{2011ApJ...733...87S} have shown an explicit dependence of $\\Ssfr$ on $\\Sigma_\\ast$.  Their ``extended Schmidt law'' --- one that incorporates the dependence on $\\Sigma_\\ast$ --- is consistent with the self-regulation model proposed by \\citet[][hereafter \\citetalias{2010ApJ...721..975O}]{2010ApJ...721..975O}, who find $\\Ssfr \\propto \\Sigma_g \\Sigma_\\ast^{1/2}$ for galaxies with a constant-scale-height stellar disk\\footnote{ In detail, the star-formation law from \\citetalias{2010ApJ...721..975O} allows for star formation in starless systems, which would be prohibited by a law with an explicit dependence on the stellar mass of a galaxy. } \\citep[see also][]{2011ApJ...743...25K, 2013ApJ...776....1K}.  In their model, the explicit dependence of the star-formation law on $\\Sigma_\\ast$ is via its contribution to the vertical gravitational field of the disk.  Therefore, we also consider the correlation between the star-formation activity of a disk and its stellar mass surface density.  Previous studies considering the relation of the two-component stability level of disks and/or stellar mass surface density to the star-formation law have lacked the kinematic data necessary to measure either of these quantities dynamically.  Instead, they have used stellar mass estimates from stellar-population-synthesis modeling, which have not been directly calibrated by dynamical mass measurements in external disk galaxies \\citep[see discussion in][hereafter \\citetalias{2010ApJ...716..198B}]{2010ApJ...716..198B}.  However, with its unparalleled stellar kinematic data in the dynamically cold regime of galaxy disks and its ancillary gas data, the DiskMass Survey \\citepalias{2010ApJ...716..198B} is well suited to studying the effect of the stellar component (via its mass surface density and stability level) on star formation in galaxy disks.  Our paper is organized as follows:  We briefly discuss the relevant observational data in Section \\ref{sec:data}.  We describe our dynamical modeling in Section \\ref{sec:genmodel}; however, a more detailed discussion of this modeling approach will be presented in a forthcoming paper.  For now, we provide a brief summary of the equations used in the dynamical model in Appendix \\ref{app:dynmodel}, and we discuss the details of our sampling of the posterior probability of the model in Appendix \\ref{app:converge}.  Our probabilistic modeling is the basis for our calculations of the disk stability parameter and stellar mass surface density.  These calculations, the stability results, and a comparison of the star-formation properties of our galaxy sample with the ``Normal Spirals'' from \\citetalias{1998ApJ...498..541K} are discussed in Section \\ref{sec:stability}.  We explore any correlations among disk stability level, stellar mass surface density, and star-formation rate in Section \\ref{sec:sfrcorr}.  Among other findings, we show that the two-component disk stability parameter is anti-correlated with the star-formation activity of the disk.  In Section \\ref{sec:scaling}, we show that this anti-correlation can be predicted by considering a closed system of empirical scaling relations. Finally, we summarize and briefly discuss our results in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{ Summary of Analyses and Empirical Findings }  In this paper, we have used data from the DiskMass Survey, briefly described in Section \\ref{sec:data}, to study the relationship between dynamical properties of galaxy disks and their star-formation activity.  Unlike other local-universe surveys, the DiskMass Survey has directly measured the stellar kinematics in a sample of galaxy disks, which are critical to dynamical calculations of both stellar surface mass density and disk stability level.  Our calculations are based on a generative --- fully probabilistic --- model of the relevant gaseous and stellar observations, adopting a simple analytical model for the disk dynamics.  We introduce our probabilistic modeling approach in Section \\ref{sec:genmodel}, we briefly outline the dynamical model in Appendix \\ref{app:dynmodel}, and we discuss our usage of a sophisticated MCMC algorithm to sample the probabilistic model in Appendix \\ref{app:converge}.  A more complete discussion of our modeling approach will be the subject of a forthcoming paper.  We calculate the star-formation rate, $\\sfr$, for each of our galaxies using measurements of the 21cm radio continuum \\citep[NVSS;][]{TPKMPhD} and the calibrations from \\citet{2001ApJ...554..803Y} (see Section \\ref{sec:qeqns}). These calculations are performed by including appropriate probability distributions for the relevant quantities in our probabilistic model (equation \\ref{eq:pfull}).  We calculate an effective star-formation-rate surface density, $\\Sigma_{\\rm SFR}\\equiv\\Ssfre$, where the surface area of the disk is determined by $R_{25}$ measurements from NED (Table \\ref{tab:data}).  Combined with effective gas mass surface densities from our probabilistic model (equation \\ref{eq:avsig}), we compare our galaxy sample to the ``Normal Spirals'' from \\citetalias{1998ApJ...498..541K} in Section \\ref{sec:KS} (see Figure \\ref{fig:ks}).  We find that the galaxies in our sample have low star-formation rates relative to the KS law, but are fully consistent with the low star-formation end of the data presented by \\citet{2008AJ....136.2846B}.  We calculate the disk stability parameter derived by \\citet{2011MNRAS.416.1191R}, $\\qrw$, which incorporates both the gaseous and stellar components (see Section \\ref{sec:qeqns}) and includes corrections for the disk thickness.  In Section \\ref{sec:result}, we find a stability parameter of $\\qrw=2.0\\pm0.9$ at 1.5 $h_R$ marginalized over all galaxies in our sample. This result is comparable to other empirical assessments from the literature. In particular \\citet{2013MNRAS.433.1389R} similarly find $\\qrw\\sim2$ for their sample; however, they also find that the stellar component most often dominates the disk stability level.  This is contrary to our results, which show that the gas-only stability parameter is lower than the star-only stability parameter ($Q_g < \\mathcal{Q}_\\ast$) for 65\\% of our sample.  These different findings are more likely because of differences in the data used and the detailed analysis methods, as described in Section \\ref{sec:result}, not because of an intrinsic difference in our galaxy samples.  A stability parameter of $\\qrw\\sim2$ is also comparable to expectations from N-body simulations; however, it should be noted that these theoretical studies typically calculate the nominal (infinitely thin, single-component, collisionless) \\citet{1964ApJ...139.1217T} criterion ($\\mathcal{Q}_\\ast$), as opposed to our calculations for a multi-component, non-zero thickness disk.  An important avenue for numerical simulations of galaxy disks is to study the stability levels of disks that include realistic gas components.  The primary goal of this paper has been to explore any dependence of the star-formation rate on two dynamical properties of our sample, $\\Sse$ and $\\qrw$ (see Section \\ref{sec:sfrcorr}).  Our primary findings are: \\begin{itemize} \\item  There is a clear correlation between $\\Sse$ and $\\Ssfre$ and a clear anti-correlation of $\\qrwmin$ with $\\Ssfre$ (Figure \\ref{fig:sigsfr}). \\item The anti-correlation between $\\qrwmin$ and $\\Ssfre$ is expected given a theoretical study by \\citet{2006ApJ...639..879L}.  However, our galaxies exhibit significantly higher stability levels than in their simulations and our data show a steeper power-law slope in the relation. \\item We find that the star-formation efficiency (SFE; $\\Ssfre/\\Sge$) is correlated with $\\Sse$ (Figure \\ref{fig:sigsfg}), which is expected both observationally \\citep{2011ApJ...733...87S} and theoretically \\citepalias{2010ApJ...721..975O}.  In detail, our data are consistent with the proportionality $\\Ssfre/\\Sge \\propto \\Sse^{1/2}$, which is a limiting behavior of the theory derived by \\citetalias{2010ApJ...721..975O}; see also \\citet{2011ApJ...743...25K, 2013ApJ...776....1K}. \\item If star formation in our galaxy sample is not strongly affected by other physical properties, the quantity $\\Ssfre \\Sge^{-1} \\Sse^{-1/2}$ should be roughly constant.  Indeed, we find that this quantity is effectively uncorrelated with the large number of physical quantities we have calculated (listed in Section \\ref{sec:sfrcorr}; see the examples of $\\Sse$ and $\\qrwmin$ in Figure \\ref{fig:sigsfgs}).  However, the scatter in the data is large.  We find an error-weighted geometric mean of $\\langle\\log(\\Ssfre \\Sge^{-1}\\Sse^{-1/2})\\rangle = -3.25\\pm0.27$ in units of $(G/{\\rm pc})^{1/2}$ and a range of $-3.7\\leq\\log(\\Ssfre \\Sge^{-1}\\Sse^{-1/2})\\leq-2.7$. \\item Although the correlation is weak, we also find an indication that the SFE is anti-correlated with $\\qrwmin$ (Figure \\ref{fig:sigsfg}).  This contradicts previous observational studies \\citep{2008AJ....136.2782L}, but is roughly consistent with the expectation provided by \\citet{2006ApJ...639..879L}. \\end{itemize}  The anti-correlation between star-formation rate and disk stability parameter from Figure \\ref{fig:sigsfr} is {\\em not} seen if we instead consider the gas-only ($Q_g$) or star-only ($\\mathcal{Q}_\\ast$) stability parameters.  In terms of the effect on star formation, this result is reasonable in that neither $Q_g$ nor $\\mathcal{Q}_\\ast$ includes the gravitational effects of the other component.  Thus, our results show the importance of considering {\\em both} components in assessing the effect of the disk stability level on star formation.  In Section \\ref{sec:scaling}, we quantitatively predict $\\qrw$ at 1.5$h_R$ for the galaxies in our sample based on a closed system of scaling relations and the following global quantities for our galaxies: $\\sfr$, $\\Upsilon_K$, $\\mu_{0,K}$, $h_R$, $R_{25}$, $\\sigma_{cg}$, and $\\alpha=\\sigma_z/\\sigma_R$.  The accuracy of the prediction depends on the details of the assumed mass distributions; however, all four approaches we discuss exhibit systematic errors of roughly 35\\% (68\\% confidence interval).  Assuming our galaxy sample is representative of the overall population \\citepalias{2010ApJ...716..198B}, our calculations demonstrate that one should {\\em expect} an anti-correlation between $\\Ssfre$ and $\\qrw$, particularly for $\\qrw^{1.5h_R} < 2$.  \\subsection{ A Physical Link Between Disk Stability Level and Star Formation? }  Our use of empirical scaling relations to predict the anti-correlation between $\\Ssfre$ and $\\qrw^{1.5h_R}$ suggests that this outcome is consistent with the physical drivers of other morphological and dynamical outcomes of late-type-galaxy evolution.  However, does the anti-correlation imply a physical link between disk stability level and star formation?  The studies of \\citet{2005ApJ...626..823L, 2006ApJ...639..879L} are particularly relevant because the star-formation in their simulations is, in fact, driven by gravitational instabilities, and our data are roughly consistent with their predictions (Figures \\ref{fig:sigsfr} and \\ref{fig:sigsfg}).  However, most of our galaxies are very stable, in the regime in which \\citet{2005ApJ...626..823L} find that it is difficult to form stars ($\\qrwmin\\gtrsim1.6$).  A comparison of Figure 10 from \\citet{2006ApJ...639..879L} with our Figure \\ref{fig:sigsfr} shows that our galaxies exhibit higher stability levels at the low star-forming end.  Part of this discrepancy is due to systematic differences in our calculation of the stability parameter (e.g., our applied corrections for disk thickness); however, our data should yield relatively large stability values even after removing these systematic differences.  Thus, just as seen via ultraviolet radiation in the very extended parts of disks \\citep{2007ApJS..173..538T}, our measurements suggest that star-formation occurs in locales with high stability levels.  One theory that addresses this phenomenon is provided by \\citetalias{2010ApJ...721..975O}.  \\citetalias{2010ApJ...721..975O} \\citep[see also][]{2011ApJ...743...25K} present a model for star formation where the interstellar medium (ISM) is divided into self-gravitating and diffuse components.  The pressure in the diffuse gas --- regulated by heating, cooling, and supernova-driven turbulence --- is assumed to be balanced by the vertical gravitational forces of the disk.  A fundamental assumption is that the surface density of gravitationally bound clouds is converted to stars over a star-formation timescale; however, this timescale ($\\sim$2 Gyr) is much longer than the free-fall time.  The proportionality $\\Ssfr/\\Sigma_g \\propto \\Sigma_\\ast^{1/2}$ is a limiting behavior of their model, assuming the star-formation rate in all galaxies is similarly affected by chemical composition, turbulence, and magnetic fields.  We have shown in Figure \\ref{fig:sigsfg} that this proportionality is consistent with our data, albeit with significant scatter.  We have not assessed the variations in metallicity, turbulent pressure, or magnetic field strength in our sample, such that this may explain some of the scatter seen in our data.  The correspondence of our data with the limiting behavior of the \\citetalias{2010ApJ...721..975O} prediction suggests that our galaxies form stars at a rate that maintains the pressure balance in the diffuse ISM.  Such star formation does not require large-scale gravitational instabilities, but it does require the vertical gravitational field, largely effected by the stellar component in most cases, to maintain the total pressure in the disk midplane. Thus, star-formation can be active in the disks of our galaxies, despite our rather large measurements of $\\qrw$.  This may argue against a physical link between $\\qrw$ and $\\Ssfre$ at large $\\qrw$.  However, the consistency of the trend seen in Figure \\ref{fig:sigsfr} with the prediction of \\citet{2006ApJ...639..879L} at low $\\qrw$ is suggestive of the role gravitational instabilities play in this small subset ($\\sim$15\\%) of our sample.  Therefore, a direct physical link between the stability level of galaxy disks and their star-formation activity may only be relevant to a small subset of the galaxies in the local universe; however, the story is most likely very different at earlier cosmic epochs \\citep[see, e.g.,][]{2009MNRAS.397L..64A}.  \\subsection{ Spiral Structure Effects }   A detailed analysis of the spiral-arm strength in our galaxy sample is beyond the scope of this paper; however, Figure 9 from \\citetalias{2010ApJ...716..198B} shows that spiral arms are easily discernible in all of our sample galaxies. From N-body simulations, we expect spiral structure to be more easily generated in disks with lower $\\mathcal{Q}_\\ast$.  If we assume that the same is true for $\\qrwmin$, we might expect effects related to spiral arms to be more evident in galaxies with lower $\\qrwmin$.  In their appendix, \\citetalias{2010ApJ...721..975O} discuss the effects of spiral structure on their equilibrium model.  They find that the azimuthally averaged (or disk-averaged) star-formation-rate surface density should be significantly lower than the true value if the contrast between the arm and inter-arm gas mass surface density is sufficiently high.  If this effect were evident in our data, we may expect the residuals of our galaxy sample about $\\Ssfre \\Sge^{-1}\\Sse^{-1/2} = 10^{-3.25}$ $(G/{\\rm pc})^{1/2}$ to be related to the strength of the spiral structure and therefore to $\\qrwmin$.  However, Figure \\ref{fig:sigsfgs} shows no such relation.  This may be because (1) $\\qrwmin$ is not a good proxy for the mass-loading of spiral-arm density waves, (2) our galaxies have all been similarly affected by spiral structure such that these properties have only changed the mean value of $\\Ssfre\\Sge^{-1}\\Sse^{-1/2}$, (3) the timescales involved in the passage or lifetime of spiral arms is such that the equilibrium is never fully realized or substantially altered within the spiral-arm regions, or (4) the scatter in $\\Ssfre \\Sge^{-1}\\Sse^{-1/2}$ caused by spiral arms or other physical processes has obscured the relation.  Some of the scatter in our measurements of $\\Ssfre\\Sge^{-1}\\Sse^{-1/2}$ can be attributed to systematic error; however, there is room for intrinsic scatter as well.  It is of great interest to understand the scatter in $\\Ssfre\\Sge^{-1}\\Sse^{-1/2}$, as it relates to spiral structure and/or other physical properties that can affect how stars form."
    },
    "1402/1402.0383_arXiv.txt": {
        "abstract": "{{\\em Fermi}/LAT observations of star-forming galaxies in the $\\sim$0.1-100\\gev\\ range have made possible a first population study. Evidence was found for a correlation between \\gray\\ luminosity and tracers of the star formation activity. Studying galactic cosmic rays (CRs) in various global conditions can yield information about their origin and transport in the interstellar medium (ISM).}{This work addresses the question of the scaling laws that can be expected for the interstellar \\gray\\ emission as a function of global galactic properties, with the goal of establishing whether the current experimental data in the GeV range can be constraining.}{I developed a 2D model for the non-thermal emissions from steady-state CR populations interacting with the ISM in star-forming galaxies. Most CR-related parameters were taken from Milky Way studies, and a large number of galaxies were then simulated with sizes from 4 to 40\\,kpc, several gas distributions, and star formation rates (SFR) covering six orders of magnitude.}{The evolution of the \\gray\\ luminosity over the 100\\kev-100\\tev\\ range is presented, with emphasis on the contribution of the different emission processes and particle populations, and on the transition between transport regimes. The model can reproduce the normalisation and trend inferred from the \\textit{Fermi}/LAT population study over most of the SFR range. This is obtained with a plain diffusion scheme, a single diffusion coefficient, and the assumption that CRs experience large-scale volume-averaged interstellar conditions. There is, however, no universal relation between high-energy \\gray\\ luminosity and star formation activity, as illustrated by the scatter introduced by different galactic global properties and the downturn in \\gray\\ emission at the low end.}{The current \\textit{Fermi}/LAT population study does not call for major modifications of the transport scheme for CRs in the Milky Way when extrapolated to other systems, probably because the uncertainties are still too large. Additional constraints may be expected from doubling the {\\em Fermi}/LAT exposure time and later from observing at TeV energies with the \\textit{Cherenkov Telescope Array}.} ",
        "introduction": "\\label{intro}  \\indent Star-forming galaxies harbour large-scale populations of energetic particles accelerated in violent phenomena associated with stellar evolution, such as supernova explosions, colliding-wind binaries, or pulsar wind nebulae. These galactic cosmic rays (hereafter CRs) interact with the various components of the interstellar medium (ISM): gas, turbulence, radiation, and magnetic fields. These interactions determine the transport of CRs from sources to intergalactic medium, and are responsible for their confinement and accumulation in the galactic volume over a large number of acceleration episodes (of the order of $\\sim10^5$ for the Milky Way). The existence of such large-scale populations is attested by various extended emissions, notably in radio through synchrotron emission, but also in \\grays\\ through inverse-Compton scattering, Bremsstrahlung, and inelastic nuclear collisions. These emissions are signatures of the combined processes of CR acceleration and transport, and their interpretation in the frame of other astronomical data can provide us with a better picture of the life cycle of CRs \\citep[see][for a review]{Strong:2007}.  \\indent Understanding CRs is relevant to many fields beyond high-energy astrophysics. This is mostly because they are more than a simple side effect of the end stages of stellar evolution. Galactic cosmic rays interact with the ISM and can significantly modify the environment in which they are flowing. They can heat and ionise the gas, especially at the centre of dense molecular clouds, thereby impacting star formation conditions \\citep{Papadopoulos:2010}; through their diffusion, CRs can alter the turbulence of the interstellar medium, with consequences on their own transport \\citep{Ptuskin:2006}; they can contribute to large-scale outflows, an important feedback effect on star formation over cosmic times \\citep{Hanasz:2013}; CRs may play a role in shaping large-scale galactic magnetic fields through a dynamo process \\citep{Hanasz:2004}; and last, CRs from classical astrophysical objects can hide the signatures of more exotic components of the Universe such as dark matter \\citep{Delahaye:2010}. For all these reasons, understanding the non-thermal content and emission of star-forming galaxies is crucial.  \\indent The current generation of high-energy (HE) and very-high-energy (VHE) \\gray\\ telescopes has enabled us to detect several external galaxies that shine at photon energies $>100$\\mev\\ as a result of CRs interacting with their ISM (and not because of a central black hole activity, as is the case for the vast majority of detected \\gray-emitting external galaxies). The {\\em Fermi}/LAT space-borne pair-creation telescope has permitted the detection of five external galaxies at GeV energies (SMC, LMC, M31, M82, and NGC253), with two other systems being possible candidates \\citep[NGC4945 and NGC1068, whose emission may be contaminated by central black hole activity; see][]{Abdo:2009l,Abdo:2010d,Abdo:2010e,Abdo:2010f,Ackermann:2012}. In the same time, the HESS and VERITAS ground-based Cherenkov telescopes have permitted the detection at TeV energies of NGC253 and M82, respectively \\citep{Acero:2009,Abramowski:2012,Acciari:2009}. Star-forming galaxies can now be considered as an emerging class of \\gray\\ sources, with a notable increase of the number of detected objects from the previous generation of instruments \\citep[only the LMC was detected by CGRO/EGRET; see][]{Sreekumar:1992}.  \\indent Even with such a limited sample, it is tempting to perform a population study to find clues about the main parameters regulating CR populations. From the {\\em Fermi}/LAT sample of five objects plus an estimate for the global emission of the Milky Way and an upper limit for M33, the work of \\citet{Abdo:2010f} showed first hints for a slightly non-linear correlation between \\gray\\ luminosity and star formation rate (SFR). A more systematic study by \\citet{Ackermann:2012}, including upper limits for about sixty luminous and ultra-luminous infrared galaxies, found further evidence for a quasi-linear correlation between \\gray\\ luminosity and tracers of the star formation activity. Such a relationship can be expected to first order, because the production of CRs is expected to scale with the rate of massive star formation; but it is surprising that the correlation is so close to linear, over such a wide range of galactic properties, from dwarfs like the SMC to interacting systems like Arp 220.  \\indent The interpretation of this correlation therefore raised the question of the scaling laws that can be expected for the diffuse \\gray\\ emission as a function of global galactic properties. Clarifying this at the theoretical level is required to establish whether the current experimental data can be constraining in any way. This is all the more needed because the sample of detected objects is small, the coverage over the GeV-TeV range is uneven, and the spectral characterization is still rather limited. In addition, the determination of global properties of external systems, such as total mass or star formation rate, is in itself a challenge, and there is most likely some intrinsic scatter in the actual physical correlations.  \\indent That star-forming galaxies might become a class of \\gray\\ sources was anticipated early, well before the launch of the \\textit{Fermi} satellite \\citep[see for instance][]{Volk:1996, Pavlidou:2001,Torres:2004}. The interest in the subject extended beyond the study of nearby objects, because this emission cumulated over cosmic time might account for a significant fraction of the isotropic diffuse \\gray\\ background \\citep{Pavlidou:2002,Thompson:2007, Ackermann:2012}. Before and after the launch of the \\textit{Fermi} satellite, the \\gray\\ emission from individual starbursts was studied by several authors, with the focus being mainly on M82 and NGC253 \\citep{Persic:2008,deCeadelPozo:2009,Lacki:2011,Paglione:2012}.  \\indent The purpose of this paper is to provide a more global picture of the evolution of the \\gray\\ interstellar emission from star-forming galaxies, in the context of the recent population study conducted on the basis of {\\em Fermi}/LAT observations. Instead of trying to fit a model to a particular object, I defined a global framework by fixing as many cosmic-ray-related parameters as possible from Milky Way studies, leaving only a few global quantities as independent variables. Here, I present a generic model for the transport of CRs and the associated non-thermal interstellar emissions in star-forming galaxies. Simulating systems with various sizes and gas distributions thus allows investigating the impact of global galactic parameters on the \\gray\\ output. In a first part, I introduce the different components and related assumptions of the model, together with the set of synthetic galaxies that was considered. Then, the \\gray\\ emission over 100\\kev-100\\tev\\ for this series of synthetic galaxies is presented, and its evolution with global properties is analysed. In a subsequent part, the scaling of \\gray\\ luminosity with global galactic properties is derived for soft, high-energy, and very-high-energy \\gray\\ ranges, and is compared with observations in the case of the high-energy band. Last, the impact of the model parameters is discussed, and the connection to the far-infrared - radio correlation is addressed. ",
        "conclusions": "\\label{conclu}  \\indent The interaction of galactic cosmic rays (CRs) with the interstellar medium (ISM) produces \\grays\\ through hadronic interactions, inverse-Compton scattering, and Bremsstrahlung. This interstellar emission has been observed in the Milky Way for decades and provided a way to indirectly probe CRs outside the solar system, especially the dominant nuclei component. The current generation of high-energy (HE) and very-high-energy (VHE) \\gray\\ instruments has enabled us to extend this kind of study to several other star-forming galaxies. Five external systems, possibly seven, were recently detected at GeV energies with the \\textit{Fermi}/LAT. Combined with upper limits on about sixty other star-forming galaxies, evidence was found for a quasi-linear correlation between \\gray\\ luminosity and tracers of the star formation activity. This raised the question of the scaling laws that can be expected for the interstellar \\gray\\ emission as a function of global galactic properties, with the aim of knowing whether this early population study can constrain the origin and transport of CRs.  \\indent In the present work, I introduced a model for the non-thermal emissions from steady-state CR populations interacting with the ISM in star-forming disk galaxies. The model is two-dimensional and includes realistic CR source distributions as well as a complete treatment of their transport in the galactic disk and halo (in the diffusion approximation). This contrasts with most previous works, which relied on a one-zone approach and typical time-scale estimates. Prescriptions were used for the gas distribution, the large-scale magnetic field, and the interstellar radiation field. Cosmic-ray-related parameters such as injection rates and spectra or diffusion properties were taken from Milky Way studies. The model is set to reproduce the average interstellar conditions of the Galaxy, its global non-thermal emission, and the far-infrared - radio correlation. A series of star-forming galaxies with different sizes, masses, and gas distributions was simulated to assess the impact of global galactic parameters on the \\gray\\ output. The link between gas content and star formation was made from the Schmidt-Kennicutt relation. The full set covered almost six orders of magnitude in star formation rate (SFR). The emission was examined over 100\\kev-100\\tev, with dedicated discussions for instrumental bands: 100\\kev-10\\mev\\ (band I, soft \\grays), 100\\mev-100\\gev\\ (band II, HE \\grays), and 100\\gev-100\\tev\\ (band III, VHE \\grays).  \\indent In band I, the emission is dominated by inverse-Compton emission, with Bremsstrahlung contributing at the $\\sim20$\\% level. With increasing average gas surface density, the luminosity rises more than linearly, with a power-law index on SFR $\\sim$1.2-1.3, depending on the gas distribution. This is the result of the growing contribution of secondary leptons. The latter eventually accounts for 70-80\\% of the emission.  \\indent In band II, the emission is dominated by pion-decay emission, with Bremsstrahlung contributing at the $\\sim20$\\% level, and inverse-Compton being equally important for small galaxies with low average gas surface density. With increasing average gas surface density, the luminosity rises more than linearly, with a power-law index on SFR initially in the range $\\sim$1.4-1.7 and then decreasing to $\\sim$1.1-1.2. This behaviour is the result of the progressive shift of the CR nuclei population from a diffusion-dominated regime towards a loss-dominated regime.  \\indent In band III, the emission is dominated by inverse-Compton emission for small galaxies with low average gas surface density. For larger and denser systems, pion decay is the main emission mechanism, accounting for more than 95\\% of the luminosity for the densest ones. With increasing average gas surface density, the luminosity rises more than linearly, with a power-law index on SFR in the range $\\sim$1.4-1.5. The behaviour is the same as in band II, with emitting particles being closer to diffusion-dominated, and with inverse-Compton becoming increasingly important at the low end.  \\indent Without any specific tuning, the model was able to account for the normalisation and trend inferred from the {\\em Fermi}/LAT population study over almost the entire range of SFRs tested here. This suggests that the \\textit{Fermi}/LAT population study does not call for major modifications of the scheme inferred for CRs in the Milky Way when extrapolated to other systems, probably because uncertainties are still too large. A good match can be obtained with a plain diffusion scheme, a single diffusion coefficient, and the assumption that CRs experience large-scale volume-averaged interstellar conditions. There is no need for a strong dependence of the diffusion properties on interstellar conditions, and no impact of strong galactic winds in the most active star-forming systems. The only requirement in the framework used here is that small galaxies with a disk radius of a few kpc have halos of at least 2\\,kpc in half height. There is, however, no universal relation between high-energy \\gray\\ luminosity and star formation activity, as illustrated by the downturn in \\gray\\ emission at low SFR values and the scatter introduced by different galactic global properties. Varying the galaxy size, gas distribution, and gas density introduces a variation in the 100\\mev-100\\gev\\ luminosity at a given SFR by up to a factor of 3, while varying the SFR for a given gas layout was found to add another 50\\% of scatter. Interestingly, the same model accounted pretty well for the far-infrared - radio correlation over most of the range of galactic properties tested here, with a remarkably small scatter of $\\leq 30$\\% at any infrared luminosity and a decrease in synchrotron emission at the low end.  \\indent Progress on this kind of population study is expected from the doubling of the {\\em Fermi}/LAT exposure, especially at low SFRs with the detection of M33 and the resolving of the Magellanic Clouds. An additional test of our understanding of interstellar \\gray\\ emission from star-forming galaxies may come from the 100\\gev-100\\tev\\ range. In this band, the scaling of luminosity with SFR has a higher non-linearity favourable to distant but more active objects, and the next-generation {\\em Cherenkov Telescope Array} with its ten-fold improvement in sensitivity is expected to perform constraining observations.  \\begin{acknowledgement} I acknowledge support from the European Community via contract ERC-StG-200911 for part of this work. I wish to thank Olivier Bern\\'e and D\\'eborah Paradis for their assistance on interstellar radiation field models, and Andy Strong for his help with the GALPROP package. I also thank Keith Bechtol and Guillaume Dubus for reading and commenting on a draft version of the manuscript. \\end{acknowledgement}"
    },
    "1402/1402.2386_arXiv.txt": {
        "abstract": "For the first time we reconstruct the magnetic helicity density of global axisymmetric field of the Sun { using} method proposed by \\citet{2003AdSpR..32.1835B} and \\citet{pip13M}. To determine the components of the vector potential, we apply the gauge which is typically employed in mean-field dynamo models. This allows for a direct comparison { of reconstructed helicity} with the predictions from the mean-field dynamo models. We apply the method to two different data sets: the synoptic maps of line-of-sight (LOS) magnetic field from the Michelson Doppler Imager (MDI) on board of Solar and Heliospheric Observatory (SOHO) and vector magnetic field measurements from Vector Spectromagnetograph (VSM) on Synoptic Optical Long-term Investigations of the Sun (SOLIS) system. Based on the analysis of MDI/SOHO data, we find that in solar cycle 23 the global magnetic field had positive (negative) magnetic helicity in the northern (southern) hemisphere. This hemispheric sign asymmetry is opposite to helicity of solar active regions, but it is in agreement with the predictions of mean-field dynamo models. { The data also suggest that the hemispheric helicity rule may have reversed its sign in early and late phases of cycle 23.} Furthermore, the data indicate an imbalance in magnetic helicity between the northern and southern hemispheres. This imbalance seem to correlate with the total level of activity in each hemisphere in cycle 23. Magnetic helicity for rising phase of cycle 24 is derived from SOLIS/VSM data, and qualitatively, its latitudinal pattern is similar to the pattern derived from SOHO/MDI data for cycle 23. ",
        "introduction": "The generation of the magnetic field in the Sun is tightly related with the convective helical motions. In the framework of axisymmetric dynamos, the magnetic field is typically decomposed into toroidal and poloidal components. \\citet{P55} suggested that solar dynamo can be represented as a periodic transformation of poloidal magnetic field, $\\bar{\\mathbf{B}}^{(p)}={\\bar{B}}_{r}\\mathbf{e}_{r}+{\\bar{B}}_{\\theta}\\mathbf{e}_{\\theta}$, into toroidal field $\\bar{\\mathbf{B}}^{(t)}={\\bar{B}}_{\\phi}\\mathbf{e}_{\\phi}$ (via the differential rotation) and the reverse transformation of $\\bar{\\mathbf{B}}^{(t)}$ to $\\bar{\\mathbf{B}}^{(p)}$ by helical convective motions. Further development of dynamo theory showed that the two processes produce helical magnetic fields on both small and large spatial scales \\citep{pouquet-al:1975a,pouquet-al:1975b}, and that the conservation of magnetic helicity is an important factor for the dynamical quenching of large-scale magnetic field generation \\citep{kleruz82,1991ApJ...376L..21C,1992ApJ...393..165V,kle-rog99,2000A&A...361L...5K,2005PhR...417....1B}.  Early observations of various proxies of magnetic/current helicity established what is now known as the hemispheric helicity rule: magnetic fields of active regions exhibit preferentially negative (positive) helicity in the northern (southern) hemisphere \\citep[][and references therein]{see1990SoPh,pev95,zetal10}.  On the other hand, some researchers \\citep[e.g.,][]{2003AdSpR..32.1835B,bl-br2003,warn2011,pip2013ApJ,pip13M} argued that magnetic helicity of large-scale (global) axisymmetric field should be positive/negative in the northern/southern hemisphere. Furthermore, the mean-field dynamo models predict reversals of the sign of helicity in association with the propagation of dynamo wave inside the convection zone \\citep[e.g.,][]{warn2011,pip13M}. Reversals of the hemispheric helicity rule have been reported in observations, but the results seems inconclusive, with some researchers reporting the presence of such reversals in early/late phases of solar cycle \\citep{bao2000,hag05}, while others are questioning their existence  \\citep{pev2001,pev2008,2013ApJ...772...52G}. {Such apparent controversy may be resolved via direct comparison of observations with the model predictions. Since dynamo models can provide a detailed information about the distribution of magnetic helicity of global axisymmetric field in the convection zone and near the photosphere, it is highly desirable to directly compare these model estimates with the observations.}  Theory suggests that magnetic helicity on small and large scales should have opposite sign \\citep[e.g.,][]{pouquet76,seehafer96,bl-br2003,2005PhR...417....1B,pip2013ApJ,pip13M}. Then, the small-scale helicity may dissipate on small spatial scales subject to the Ohmic dissipation \\citep[e.g.,][]{pouquet76} or the helicity of both signs could emerge through the solar photosphere. Early measurements of vertical component of large-scale current helicity density, \\citep[e.g.,][]{2000ApJ...528..999P,2010ApJ...720..632W} found that in its sign, the large-scale magnetic fields follow the same hemispheric helicity rule as the active regions. These early studies concentrated on spatial scales larger than the active regions but smaller then the solar hemisphere. One should note that in the framework of mean-field dynamo, the magnetic fields of active regions represent the ``small-scale'' fields, while the large-scale fields refer to spatial scales comparable with the size of solar hemisphere. In this article, we address the helicity determination for large-scale magnetic fields as defined by mean-field dynamo theory.  To avoid confusion with previous studies, we use the terms ``global'' and ``large-scale'' to refer to magnetic fields on spatial scales much larger then active regions. We reconstruct the magnetic helicity of global axisymmetric field using the approach suggested by \\citet{2003AdSpR..32.1835B} and  \\citet[][see Section 2 below]{pip13M}. Section 3 describes the data sets and the reduction procedure. Section 4 presents our main results, and in Section 5 we discuss our findings. ",
        "conclusions": "Using synoptic charts from SOHO/MDI, we, for the first time, reconstruct the magnetic helicity density of global axisymmetric field of the Sun. In solar cycle 23 the global axisymmetric magnetic field exhibits positive magnetic helicity in the northern hemisphere, and negative one in the southern. In general, such reconstructions require a knowledge of the $\\bar{B}_{r}$ and $\\bar{B}_{\\phi}$ components of the axisymmetric magnetic field. In the past, vector global magnetic field components were reconstructed via various approaches \\citep{2000ApJ...528..999P,2005ApJ...620L.123U,hoeks10,2012SoPh..280..379M}. Here we used synoptic charts of LOS magnetic field corresponding to different longitudinal offsets relative to central meridian to compute $\\bar{B}_{r}$ and $\\bar{B}_\\phi$, see equation (\\ref{eq:torf}). Derived $\\bar{B}_{r}$ agrees well with the $\\bar{B}_{r}$ provided by the MDI team, which we see as indirect validation of our method.  Based on the analysis of dynamo equations, \\citet{2003AdSpR..32.1835B} suggested that the magnetic helicity of global magnetic field in solar cycle 23 should be positive/negative in the northern/southern hemisphere. \\citet{pip13M} analyzed the distributions of magnetic helicity for large- and small-scale magnetic fields in the axisymmetric mean-field dynamo taking into account the conservation of the total magnetic helicity in the dynamo processes. They concluded (see their Figures 2e,d and 5) that magnetic helicity density of large-scale field should have positive sign in the northern hemisphere and negative sign in the southern hemisphere during the most part of the magnetic cycle. Our present results provide observational support to these early theoretical predictions. We find that during most of cycle 23, global magnetic fields exhibited a persistent pattern of positive/negative helicity in the northern/southern hemispheres.  In respect to helicity of active region magnetic fields (small-scale in the framework of this discussion), the hemispheric helicity rule is negative/positive in the northern/southern hemispheres \\citep[][and references therein]{see1990SoPh,pev95,zetal10}. Taken together, these two results support the notion that the solar dynamo creates helicity of two opposite signs as was suggested in early papers. However, helicity of both signs seem to cross the solar photosphere. We further found that the hemispheric helicity rule for global magnetic fields exhibits sign-reversals in early and late phases of cycle 23. If the helicities of small- and large-scale fields are tied together, this should imply a need for similar reversals in the hemispheric helicity rule for active regions. Alas, while some researchers claimed observing reversals in the hemispheric helicity rule near the minimum of solar cycle 22 and 23 \\citep{bao2000,hag05}, others were not able to find them \\citep{pev2001,pev2008,2013ApJ...772...52G}. Clearly, this question about possible reversals of the hemispheric helicity rule for active region magnetic fields needs to be re-examined. Although, the predictions of the model by \\cite{pip13M} agree qualitatively with the results reported in our paper,  there are some differences related to the shape of the helicity density patterns. For example, our present results suggest that magnetic helicity density pattern of the same sign can extend from the equator to the poles which is not seen in the model.  Similarly, the pattern of $\\bar{\\mathbf{A}}\\cdot\\bar{\\mathbf{B}}$ of reverse sign penetrates to equatorial regions during the minima of cycle around year 1997 and 2009. If the reversals of the hemispheric helicity rule are real, this will pose a challenge for some proposed mechanisms of helicity generation \\citep[e.g., helicity  generation by the differential rotation,][]{2000JGR...10510481B}.  { As an alternative explanation, our results could be interpreted in the framework of helicity of axisymmetric and non-axisymmetric parts of global magnetic fields. In that model, the axisymmetric component of helicity (derived in this article) follows the hemispheric helicity rule of opposite in sign to the non-axisymmetric component (associated with active regions).  Such a possibility was raised by Zhang (2006), who used Berger \\& Ruzmaikin (2000) data to show that helicity flux of non-axisymmetric modes was opposite in sign to helicity of axisymmetric (m=0) mode.}    { One may also question the importance of vector magnetic field measurements for studies of global helicity, when the line-of-sight data seem to provide reasonable results. Here we presented the first ever derivations of magnetic helicity density of global field based on vector synoptic maps. While we see similarities in distribution of global helicity derived from LOS and vector data, there are also some differences. For example, synoptic maps of toroidal field derived from LOS data show more or less uniform distribution of the magnetic field polarity of one sign suggested by the the Hale polarity law. In addition to that pattern, vector field maps show mix of two different polarities: one is more concentrated and other is somewhat diffused (Figure 3b).  The diffuse component (of toroidal field) seems to correspond to a trailing polarity field. Such component of large-scale field is not present in the maps of toroidal flux derived from the LOS  magnetic fields. The vector field data are limited to cycle 24, and thus, are insufficient to conclude if there are any changes in this diffuse component of toroidal field with solar cycle. This and other differences between derivations based on LOS or vector field data require further investigation.}  Our findings indicate that helicity of the large-scale magnetic fields is imbalanced between the northern and the southern hemispheres in different phases of solar cycle. However, when taken over the entire cycle, the positive and negative helicity of large-scale magnetic field is well-balanced. Indirectly, this is in agreement with \\citet{geor2009}, who found that helicity injection through the solar photosphere associated with active region magnetic fields is well-balanced over the solar cycle 23. On the other hand, \\citet{2012ApJ...758...61Y} reported significant imbalance between helicity fluxes of northern and southern hemispheres. Our findings (Figure 8c) allow to reconcile \\citet{geor2009} and \\citet{2012ApJ...758...61Y} conclusions.  {Due to limitations of the existing datasets, the observational studies of helicity often refer to proxies of current helicity density. It is  usually assumed that these proxies represent magnetic helicity sufficiently well. Contrary to that, we find that while the general tendencies are similar in magnetic and current helicity densities, there are differences, for example, in small-scale patterns, which may be present in one helicity proxy but are absent in the other. For example,} proxy of current helicity, $\\bar{B}_{r}\\left(\\nabla\\times\\bar{\\mathbf{B}}\\right)_{r}$, exhibits a distinct ``zebra'' pattern, but no such pattern is present in the distribution of magnetic helicity. Early, \\citet{2000ApJ...528..999P} and \\citet{2013ApJ...772...52G} reported similar pattern in current helicity density of large-scale magnetic fields. The pattern could also be expected from spatial structure of the dynamo wave of the large-scale magnetic field components $\\bar{B}_{r}$ and $\\bar{B}_{\\phi}$, which are illustrated in Figure \\ref{fig:Components-of-the}a. We note that in the equatorial regions the inequality $\\bar{B}_{r}\\bar{B}_{\\phi}<0$ holds for the most part of the sunspot cycle (see, Figure \\ref{fig:Components-of-the}a). We also found that modes $b_{r}^{(3)}$ and $b_{\\phi}^{(2)}$ dominate, which means that $\\bar{B_{r}}\\left(\\boldsymbol{\\nabla}\\times\\mathbf{\\bar{B}}\\right)_{r}\\sim b_{\\phi}^{(2)}b_{r}^{(3)}P_{2}P_{3}$, here the sign of the $b_{\\phi}^{(2)}b_{r}^{(3)}$ defines the hemispheric sign rule and the product $P_{2}P_{3}$ defines that zebra pattern as illustrated in Figure \\ref{fig:tot}c and Figure \\ref{fig:tot-1}. Thus, the results shown in Figure \\ref{fig:tot}c are expected for any dynamo model that qualitatively reproduces Figure \\ref{fig:Components-of-the}a.  Finally, keeping in mind the approximations which were used in the reconstruction of components of the global magnetic field of the Sun, our results should be considered as preliminary. Further development in this direction is likely to shed more light on the role of magnetic helicity in global solar and astrophysical dynamos."
    },
    "1402/1402.5429_arXiv.txt": {
        "abstract": "\\noindent An intriguing discrepancy emerging in the concordance model of cosmology is the tension between the locally measured value of the Hubble rate, and the `global' value inferred from the cosmic microwave background (CMB). This could be due to systematic uncertainties when measuring $H_0$ locally, or it could be that we live in a highly unlikely Hubble bubble, or other exotic scenarios. We point out that the global $H_0$ can be found by extrapolating $H(z)$ data points at high-$z$ down to $z=0$. By doing this in a Bayesian non-parametric way we can find a model-independent value for $H_0$. We apply this to 19 measurements based on differential age of passively evolving galaxies as cosmic chronometers. Using Gaussian processes, we find $H_0=64.9 \\pm 4.2$ km s$^{-1}$ Mpc$^{-1}$ $(1\\sigma)$, in agreement with the CMB value, but reinforcing the tension with the local value. An analysis of possible sources of systematic errors shows that the stellar population synthesis model adopted may change the results significantly, being the main concern for subsequent studies. Forecasts for future data show that distant $H(z)$ measurements can be a robust method to determine $H_0$, where a focus in precision and a careful assessment of systematic errors are required. ",
        "introduction": "There is a strong tension, recently quantified by \\cite{verde2013}, between the value of the Hubble constant $H_0$ derived by {\\it Planck} \\citep{planck} from anisotropies in the cosmic microwave background (CMB): $67.3 \\pm 1.2$ km s$^{-1}$ Mpc$^{-1}$, and the value from local measurements: $73.8 \\pm 2.4$ km s$^{-1}$ Mpc$^{-1}$ \\citep{H0riess}. While the latter measurement is based on local measurements, the former infers a global value for the Hubble constant within a cosmological model.  There remains disagreement about the local value of $H_0$ depending on the distance indicator used to measure it, which hints the discrepancy with {\\it Planck} could be the result of systematic errors. \\cite{H0riess} calibrated the SNe Ia distances with three indicators: distance to NGC 4258 based on a megamaser measurement, parallax measurements to Milk Way cepheids (MWC) and cepheids observations and a revised distance to the Large Magellanic Cloud (LMC). Contrarily, calibrating the SNe Ia with the tip of red-giant branch, \\cite{tammann} provides $H_0$ $=$ 63.7 $\\pm$ 2.3 km s$^{-1}$ Mpc$^{-1}$. This shows how crucial is the first-step calibration in the distance ladder to measure $H_0$.  However, there are several local $H_0$ measurements with higher values. \\cite{riess2012} found $H_0$ $=$ 75.4 $\\pm$ 2.9 km s$^{-1}$ Mpc$^{-1}$ by using cepheids in M31. With a mid-infrared calibration for the cepheids, \\cite{freedman2012} derived $H_0$ $=$ 74.3 $\\pm$ 2.1 km s$^{-1}$ Mpc$^{-1}$, and with 8 new classical cepheids observed in galaxies hosting SNe Ia \\cite{fiorentino2013} got $H_0$ $=$ 76.0 $\\pm$ 1.9 km s$^{-1}$ Mpc$^{-1}$. By using HII regions and HII galaxies as distance indicators, \\cite{chavez2012} obtained $H_0$ $=$ 74.3 $\\pm$ 3.1(random) $\\pm$ 2.9 (syst.) km s$^{-1}$ Mpc$^{-1}$.  Some of these are over 4$\\sigma$ away from the CMB-derived value. See Fig. \\ref{figH0} for a plot of different measurements of $H_0$.  \\begin{figure} \\centerline{\\epsfig{figure=graf_values_H0.eps,width=0.5\\textwidth} \\hskip 0.1in} \\caption{{\\it Different measurements of $H_0$.} The figure shows how the result obtained in this work (GaPP) is compared to other determinations of $H_0$. The points refer to the following references: Planck \\citep{planck}, TR \\citep{tammann}, Efs1 \\citep{efstathiou} with one anchor, Efs3 with three anchors, R11 \\citep{H0riess}, R12 \\citep{riess2012}, Freed \\citep{freedman2012}, Fior \\citep{fiorentino2013} and Chavez \\citep{chavez2012}.} \\label{figH0} \\end{figure}  A variety of different physical effects could explain such a discrepancy. It could just be cosmic variance: as we can observe the Universe from only one position, we are not able to realize the global parameters from the local parameters, as in the local expansion rate for instance. If we live in an locally underdense region, a ``Hubble bubble'', a higher value for $H_0$ is obtained compared to the global value. This effect was carefully addressed by \\cite{cosmic_variance} through a modelling of the statistics of matter distribution which provides the distribution of the gravitational potential at the observer. The outcome is that cosmic variance can alleviate the tension, but a complete elimination requires a very rare fluctuation \\citep{cosmic_variance,cosmic_variance2}.  Another way to look at the problem is to consider that the discrepancy may indicate new physics, such as massive neutrinos \\citep{hu}, or alternative dark energy models \\citep{cde,xcdm}.  Recently, some analyses were performed trying to identify sources of systematic errors in order to remove or alleviate the tension. For example, by using only the geometric maser distance to NGC 4258 of \\cite{humphreys} as an anchor, \\cite{efstathiou} revisited \\cite{H0riess} analysis and derived $H_0=70.6 \\pm 3.3$ km s$^{-1}$ Mpc$^{-1}$, while combining with LMC and MWC anchors the value is $72.5 \\pm 2.5$ km s$^{-1}$ Mpc$^{-1}$, alleviating the tension. The {\\it Planck} data were also reanalysed by \\cite{spergel}, where it was claimed that the 217 GHz $\\times$ 217 GHz detector is responsible for some part of the tension. Their new Hubble constant without the 217 GHz $\\times$ 217 GHz detector is slightly higher: $H_0 = 68.0 \\pm 1.1$ km s$^{-1}$ Mpc$^{-1}$.  With so many alternatives, progress can be achieved by developing new ways to address the issue. We point out here that $H(z)$ data which are not calibrated on a $H_0$ estimate can be extrapolated to z=0 to provide an independent measurement of the global $H_0$. Here, the Hubble function is reconstructed in order to derive $H_0$ from 19 $H(z)$ measurements of passively evolving galaxies as cosmic chronometers \\citep{jimloeb}. Many of these are at relatively moderate and high redshifts so intrinsically probe the global value for $H_0$ rather than the local one. We use Gaussian Processes (GP), which is a non-parametric method, to obtain the value of the Hubble constant in a completely cosmological model-independent way, which is in principle not affected by the local systematics. We show the value of the Hubble constant derived in this way is lower than the standard local measurements. We obtain $H_0$ $=$ $64.9 \\pm 4.2$ km s$^{-1}$ Mpc$^{-1}$ $(1\\sigma)$, in agreement with the CMB-inferred value. A better understanding of systematic errors, especially the adopted stellar population synthesis model, is required: we show that to improve this result a big effort is necessary to decrease the errors substantially in future, and a focus on precision is worthier than the number of data.  The paper is organized as follows: in Sec. \\ref{meth} we describe GP as well as standard parametric methods adopted to constrain $H_0$. In Sec. \\ref{const} the bounds derived for the Hubble constant are displayed, followed by forecasts of constraints in Sec. \\ref{fore}. We finish the paper in Sec. \\ref{conc} with the conclusions. ",
        "conclusions": "\\label{conc}  We have applied GaPP, a non-parametric smoothing method based on Gaussian Processes, to 19 $H(z)$ measurements in order to constrain the Hubble constant $H_0$. This method does not rely on a cosmological model, so its results can be used to infer the impact of systematic errors as well as the underlying cosmological framework. We have obtained $H_0$ to be $64.9 \\pm 4.2$ km s$^{-1}$ Mpc$^{-1}$ $(1\\sigma)$, a value which is in agreement with {\\it Planck}, but in disagreement with local measurements. This supports the notion that either there are unidentified systematic errors in the local $H_0$ data, or the local value is indeed different from the global value. A better comprehension of systematic errors, especially a thorough analysis of the impact of SPS models, can improve the robustness of our results. Simulations have shown that improvements in distant $H(z)$ measurements can help pin down the global value of $H_0$."
    },
    "1402/1402.0510_arXiv.txt": {
        "abstract": "We present results from near-infrared spectroscopy of 26 emission-line galaxies at $z\\sim2.2$ and $z\\sim1.5$ obtained with the Folded-port InfraRed Echellette (FIRE) spectrometer on the 6.5-meter Magellan Baade telescope. The sample was selected from the WFC3 Infrared Spectroscopic Parallels (WISP) survey, which uses the near-infrared grism of the Hubble Space Telescope Wide Field Camera 3 to detect emission-line galaxies over $0.3\\lesssim z \\lesssim2.3$. Our FIRE follow-up spectroscopy (R$\\sim$5000) over 1.0--2.5~$\\mu$m permits detailed measurements of physical properties of the $z\\sim2$ emission-line galaxies. Dust-corrected star formation rates for the sample range from $\\sim$5--100~$\\mathrm{M}_{\\odot}$~yr$^{-1}$ with a mean of 29~$\\mathrm{M}_{\\odot}$~yr$^{-1}$. We derive a median metallicity for the sample of 12~+~log(O/H)~=~8.34 or $\\sim$0.45~Z$_{\\odot}$. The estimated stellar masses range from $\\sim$$10^{8.5}-10^{9.5}$~M$_{\\odot}$, and a clear positive correlation between metallicity and stellar mass is observed. The average ionization parameter measured for the sample, log~$U\\approx-2.5$, is significantly higher than what is found for most star-forming galaxies in the local universe, but similar to the values found for other star-forming galaxies at high redshift. We derive composite spectra from the FIRE sample, from which we infer typical nebular electron densities of $\\sim$100-400~cm$^{-3}$. Based on the location of the galaxies and composite spectra on BPT diagrams, we do not find evidence for significant AGN activity in the sample.  Most of the galaxies as well as the composites are offset in the BPT diagram toward higher [\\oiii]/\\hb\\ at a given [\\nii]/\\ha, in agreement with other observations of $z\\gtrsim1$ star-forming galaxies, but composite spectra derived from the sample do not show an appreciable offset from the local star-forming sequence on the [\\oiii]/\\hb\\ versus  [\\sii]/\\ha\\ diagram. We infer a high nitrogen-to-oxygen abundance ratio from the composite spectrum, which may contribute to the offset of the high-redshift galaxies from the local star-forming sequence in the  [\\oiii]/\\hb\\ versus [\\nii]/\\ha\\ diagram. We speculate that the elevated nitrogen abundance could result from substantial numbers of Wolf-Rayet stars in starbursting galaxies at $z\\sim2$. ",
        "introduction": "The rest-frame optical spectra of star-forming galaxies at all redshifts exhibit emission lines from which detailed physical properties can be inferred. For galaxies at the peak of cosmic star formation at $z\\sim2$, these emission lines are shifted into the near-infrared, which, combined with the intrinsic faintness of the sources, makes them difficult to observe from the ground. For this reason, relatively few near-infrared spectra of galaxies at $z\\sim2$ that cover all of the important rest-frame optical emission lines have been published to date (e.g., \\citealp{Erb06, Hainline09, Erb10, Rigby11, Belli13}).  The available near-infrared spectra of star-forming galaxies at $z\\sim2$ have revealed differences in comparison with counterparts in the local universe \\citep{Liu08, Newman13}. For example, star-forming galaxies at $z\\sim2$ tend to have higher [\\oiii]/\\hb\\ ratios at a given [\\nii]/\\ha\\ ratio than local star-forming galaxies. This observation has been attributed to more extreme interstellar medium (ISM) conditions, on average, in galaxies at high redshift, possibly as a result of higher nebular electron densities, harder ionizing radiation fields, different gas volume filling factors, or some combination of these  \\citep{Shapley05, Brinchmann08, Shirazi13, Kewley13a,  Kewley13}. The clumpy morphology and relatively high velocity dispersions observed in many of these sources \\citep{Pettini01, Forster06, Genzel08, Law09} may support the conjecture that star-formation in the early universe generally occurs in denser and higher pressure environments than those found in local star-forming galaxies. Significant contribution of active galactic nuclei (AGN) to emission line fluxes for $z\\sim2$ galaxies has also been suggested as a source of the elevated line ratios \\citep{Trump11,Trump13}.  The slitless grism spectroscopy provided by the Wide Field Camera 3 (WFC3) on the Hubble Space Telescope (HST) has enabled the discovery of large numbers of star-forming galaxies near the peak of cosmic star formation \\citep{Atek10, Atek11, Straughn11, Vanderwel11, Trump11, Brammer12}. Grism surveys such as the WFC3 Infrared Spectroscopic Parallels (WISP) survey \\citep{Atek10} are well-suited to finding low-mass star-forming galaxies at intermediate redshifts through their optical emission lines. While WFC3 grism spectroscopy detects large numbers of emission-line galaxies, it is not ideal for extracting the physical information encoded in their optical spectra. The spectral resolution ($R\\sim130$ in G141 and $R\\sim210$ in G102) is insufficient to resolve \\ha\\ from [\\nii]$\\lambda$6548, 6583, or detect line broadening due to AGN activity. Moreover, despite the broad wavelength coverage ($\\sim$0.8--1.7~$\\mu$m) of the grism, \\ha\\ is not detected for galaxies at $z\\gtrsim1.6$ and important metallicity diagnostics such as $R_{23}$~$\\equiv$~([\\oiii]$\\lambda4959,5007$+[\\oii]$\\lambda$3727)/H$\\beta$ are often inaccessible. For these reasons, ground-based spectroscopy with the new generation of near-infrared spectrometers is required to constrain the physical properties of these galaxies.  Here we present rest-frame optical spectroscopy of 26 emission-line galaxies from WISP obtained with the Folded-port Infrared Echellette (FIRE, \\citealp{Simcoe08, Simcoe10}) on the Magellan Baade 6.5 meter telescope. The sample consists of 13 sources at $z\\sim2.2$ and 13 at $z\\sim1.5$. Galaxies in the sample were selected from the WISP survey based on the detection of strong [\\oiii]$\\lambda5007$ emission (at $z\\sim2.2$) or \\ha\\ (at $z\\sim1.5$) in the G141 grism data. Follow-up near-IR spectroscopy with FIRE enables us to: (1) detect \\ha\\ for sources at $z>1.6$ and split \\ha\\ and [\\nii] in order to determine star formation rates (SFRs), (2) get accurate dust reddening estimates using the Balmer decrement, (3) infer metallicities and ionization parameters from strong lines, (4) measure diagnostic line ratios to test for nuclear activity, and (5) resolve emission line velocity dispersions. We also construct composite spectra from our sample, which allow us to investigate the statistical properties of strongly star-forming galaxies at $z\\sim2.2$ and $z\\sim1.5$ in greater detail.  This paper is structured as follows. In \\S2 we provide an overview of the WISP survey and the emission-line galaxies selected for follow-up spectroscopy. In \\S3 we discuss the Magellan FIRE spectroscopy and data reduction. In \\S4 we present the physical properties measured from the FIRE spectra, including dust obscuration, metallicity, ionization parameter, star formation rate and kinematics. In \\S5 we discuss the composite spectrum derived from the sample. In \\S6 we analyze the results, exploring the implications of our findings for the nature of starbursting galaxies at $z\\sim2$. In \\S7 we conclude with a discussion. We adopt a cosmology with $\\Omega_{M}=0.3$,  $\\Omega_{\\Lambda}=0.7$ and $H_{0}=70~\\mathrm{km~s}^{-1}~\\mathrm{Mpc}^{-1}$. ",
        "conclusions": "We have presented rest-frame optical spectra taken with Magellan FIRE of a sample of  26 emission-line galaxies at $z\\sim1.5$ and $z\\sim2.2$ selected from the HST-WFC3 grism spectroscopy of the WISP survey. The FIRE spectra provide significant additional information about these galaxies, allowing us to measure metal abundances, dust reddening, kinematics, star formation rates and important diagnostic line ratios. We also created composite spectra that reveal further details about the average properties of the sample.  Direct WFC3 imaging shows the sources to be compact, with many ($\\gtrsim$40\\%) exhibiting clumpiness or asymmetry. The galaxies are low mass ($\\sim$10$^{8.5}$--10$^{9.5}$~M$_{\\odot}$), with a median metallicity of 0.45~Z$_{\\odot}$. A clear mass-metallicity relation is found for the sample. As seen in other high-redshift samples, the ionization parameters of the galaxies in our sample are significantly higher, on average, than those found for typical star-forming galaxies in the local universe. The velocity dispersions of the emission lines for the sample range from $\\sim$50--200~km~s$^{-1}$ but usually are around 75~km~s$^{-1}$. Assuming that the dispersions are almost entirely due to gravitational motions (which is probably not the case), this implies that stars make up roughly 30\\% of the dynamical mass of the galaxies.  The composite spectra we generated from the FIRE sample reveal more detail about the average properties of emission-line galaxies at $z\\sim2.2$ and $z\\sim1.5$.  From the overall composite we infer an average electron density for the star-forming regions of 100-400~cm$^{-3}$ and put a limit on the average temperature in these regions of $T_{e}<16,800$~K from the non-detection of [\\oiii]$\\lambda4363$. The locations of galaxies and composites on the BPT diagnostic diagrams favor a starburst classification, and we find little evidence for substantial AGN contribution.  Our sample includes galaxies with unusually high [\\oiii]/\\ha\\ and [\\oiii]/\\hb\\ line ratios that are not reproduced with existing photoionization models. Composite spectra sorted on the ratio [\\oiii]/\\ha\\ indicate that those sources with high [\\oiii]/\\ha\\ ratios show stronger [\\oiii]$\\lambda$5007 and [\\neiii]$\\lambda$3869 emission relative [\\oii]$\\lambda$3727, indicating higher ionization parameters. The high [\\oiii]/(\\ha,\\hb) sources tend to be lower metallicity, which seems to be the main factor driving their higher ionization parameters. Reproducing the line ratios observed in such objects may require modifications to photoionization models to better account for the star-forming conditions in high redshift galaxies.  The well-known offset of high redshift star-forming galaxies from the local star-forming sequence in the [\\oiii]/\\hb\\ versus [\\nii]/\\ha\\ diagnostic diagram is observed in our sample, but we do not find a similar offset in the [\\oiii]/\\hb\\ versus [\\sii]/\\ha\\ diagram. We therefore investigated the possibility that the offset of high-redshift galaxies in the [\\oiii]/\\hb\\ versus [\\nii]/\\ha\\ diagram may result, at least in part, from elevated nitrogen abundances. The composite FIRE spectrum was used to infer the N/O abundance ratio, which we find to be $\\sim$0.4 dex higher in comparison with local galaxies of similar metallicity. We speculate that an elevated nitrogen abundance in high redshift star-forming galaxies may be common, and could be explained by the presence of a substantial population of Wolf-Rayet stars embedded in super star clusters in such galaxies."
    },
    "1402/1402.5932_arXiv.txt": {
        "abstract": "Astronomy produces extremely large data sets from ground-based telescopes, space missions, and simulation. The volume and complexity of these rich data sets require new approaches and advanced tools to understand the information contained therein. No one can load this data on their own computer, most cannot even keep it at their institution, and worse, no platform exists that allows one to evaluate their models across the whole of the data. Simply having an extremely large volume of data available in one place is not sufficient; one must be able to make valid, rigorous, scientific comparisons across very different data sets from very different instrumentation. We propose a framework to directly address this which has the following components: a model-based computational platform, streamlined access to large volumes of data, and an educational and social platform for both researchers and the public. ",
        "introduction": "The era of big data is upon us. Existing astronomical surveys like the Sloan Digital Sky Survey (SDSS), with a data volume of $\\sim$15 TB, already stretch the limits of what our analysis methods and software are capable of. Next-generation surveys, such as the Large Synoptic Survey Telescope (LSST), will generate an SDSS every night \\citep{2008arXiv0805.2366I}. How is a research astronomer to cope? The mishmash of mismatched survey outputs from data sets documented to varying degrees with differing catalog conventions, clunky web forms, small departmental computing clusters, and roll-your-own analysis software already wastes enormous quantities of time that could be better spent on actual science: building and testing models, looking for unexplained features, and using data visualization to build physical intuition and motivate new understanding.  We present here a new framework that addresses the major challenges of doing astronomy with big data. The framework can be roughly described as having three primary components: a model-based computational platform, streamlined access to large volumes of data, and an educational and social platform for both researchers and the public. It will let theorists compare physical models to data from a wide variety of astronomical surveys without having to understand for each survey the point-spread functions, photometric calibration, or the effect of aperture size choice on photometry. It will allow researchers to write code to run on data as if it were local, without the hassle of managing hundreds of terabytes of local data volume. Our framework will attach quantitative models and their associated fit likelihoods to astronomical objects, letting researchers ask questions like ``How many stars have colors or spectra that look like my model?''~or ``What things in the universe are unexplained by existing models?''~with straightforward database queries.  Associating quantitative models, their constituent code, and their detailed descriptions with these objects creates a mechanism to search for and compare with the relevant literature, and it provides a path into the subject for educators and the general public.  We will start by identifying current problems in data analysis, offer a high-level overview of the proposed framework, then describe the implementation of the framework in greater detail. Astronomical data sets and models are used as examples to present the ideas in a more concrete manner as this is the primary domain of the authors, but nothing presented here need be specific to astronomy. This framework as offered can be extended such that the data sets and models of many fields may overlap in new and novel ways. While pieces of this framework, such as designing large databases, data visualization, etc.~are certainly being developed in fields such as astrophysics, computer science, and human interface design, the approach described here fully integrates these and other ideas to enable new and novel scientific research and education. ",
        "conclusions": "We have presented here a framework designed to enable a new way to do science, taking advantage of the large volumes of data now available (and to come) in a way that hasn't been possible before, going beyond existing interfaces to data. The core concepts are the creation and integration of an astronomical data repository, the cross referencing of this data that understands and incorporates instrumental effects, transparent and simple programmatic access to the data, the ability to apply an arbitrary theoretical model on a computing platform, and an integrated education tool for researchers and the public alike. Managing access to increasingly larger data sets is necessary to enable us to ask the important questions. Do we understand the data? How accurate are our models? Do our models hold up with new data? What is in our data that does not fit any of our models? This last question is one of the most compelling that we look forward to addressing with this framework."
    },
    "1402/1402.4808_arXiv.txt": {
        "abstract": "A recent cross-correlation between the SDSS DR7 White Dwarf Catalog with the Wide-Field Infrared Survey Explorer (\\emph{WISE}) all-sky photometry at 3.4, 4.6, 12, and 22 microns performed by \\cite{debes11} resulted in the discovery of 52 candidate dusty white dwarfs (WDs). The 6$\\arcsec$ \\emph{WISE} beam allows for the possibility that many of the excesses exhibited by these WDs may be due to contamination from a nearby source, however. We present MMT$+$SWIRC $J$- and $H$-band imaging observations (0.5-1.5$\\arcsec$ PSF) of 16 of these candidate dusty WDs and confirm that four have spectral energy distributions (SEDs) consistent with a dusty disk and are not accompanied by a nearby source contaminant. The remaining 12 WDs have contaminated \\emph{WISE} photometry and SEDs inconsistent with a dusty disk when the contaminating sources are not included in the photometry measurements. We find the frequency of disks around single WDs in the \\emph{WISE} $\\cap$ SDSS sample to be 2.6-4.1\\%. One of the four new dusty WDs has a mass of $1.04~M_{\\odot}$ (progenitor mass $5.4~M_{\\odot}$) and its discovery offers the first confirmation that massive WDs (and their massive progenitor stars) host planetary systems. ",
        "introduction": "A new class of dusty white dwarfs (WDs) is rapidly being populated through infrared excess searches using \\emph{Spitzer}, the Wide-Field Infrared Survey Explorer (\\emph{WISE}; \\citealt{wright10}), and various ground based telescopes such as the IRTF and Gemini (\\citealt{zuckerman87,kilic05,becklin05,vonhippel07,jura07a,farihi10,debes11,barber12,xu12}). These WDs exhibit excess emission in the near- and mid-infrared due to the thermal reprocessing of light by a disk of circumstellar dust. These dust disks are the debris resultant from the tidal disruption of an asteroid that has veered off its orbit and passed within the WD's Roche lobe (\\citealt{debes02, jura03,jura08,bonsor11,debes12}). The origin of these asteroids is dynamical instability, initiated by post-main sequence stellar mass loss, whence a massive planet (if present) will start to gravitationally interact with smaller planetary bodies. After numerous interactions, an asteroid's orbit will become increasingly eccentric and in some cases approach close enough to the WD to be ripped apart by tidal forces. The debris produced by this disruption embodies a circular disk geometry after many subsequent orbits (\\citealt{debes12}).  Dust persists around a WD inside the tidal radius of the WD and outside the radius at which the equilibrium temperature is such that the dust will sublimate. Viscous torques cause the sublimated dust to accrete onto WD's surface (\\citealt{rafikov11a,rafikov11b,veras13}). This accretion results in a spectroscopically detectable pollution of the otherwise pristine WD atmosphere. Photospheric abundance analyses of these WDs show that the accreted metals originate from tidally disrupted minor bodies similar in composition to that of bulk Earth (\\citealt{zuckerman07,klein10,klein11,dufour10,dufour12,xu14}). Since at least one planet is required to perturb minor bodies out of their stable orbits (\\citealt{debes12}), photospheric pollution as well as circumstellar debris disks serve as tracers for remnant planetary systems at WDs.  Dusty disks surrounding WDs are generally assumed to be geometrically flat and optically thick (\\citealt{jura03,rafikov11a,rafikov11b}). Based on flat disk models, they vary in width from a narrow ring, a few tenths of a solar radius, to a disk filling the entire region interior to the tidal radius of the star and exterior to the radius of dust sublimation. The WDs known to host debris disks typically range in temperature from 9500-24,000 $K$ and in mass from 0.5-0.9 $M_{\\odot}$. With this work, we expand the parameter space occupied by WD+disk systems with the confirmation of the first disk orbiting a WD hotter than 24,000 $K$ and the first disk orbiting a WD more massive than 1 $M_{\\odot}$.  We present follow up near-infrared (NIR) $J$- and $H$-band photometry using the 6.5 m MMT with SWIRC of 16 WDs exhibiting a mid-infrared excess indicative of a debris disk detected by \\emph{WISE} (\\citealt{debes11}). We use the higher spatial resolution of SWIRC to confirm the presence of a debris disk or to identify a photometric contamination in the six arcsecond \\emph{WISE} beam. We start in Section \\ref{target} with the selection of our targets. Section \\ref{obs} is a description of the observations and data reduction as well as the debris disk modeling of the detected excesses. Finally, we present our results in Section \\ref{results}. ",
        "conclusions": "We have confirmed the presence of dust surrounding four WDs with excess flux detected by \\emph{WISE}. These four new dusty WDs, including the two bonafide UKIDSS$+$WISE disks, increases the total number of confirmed WD$+$disk systems from 29 to 35 (\\citealt{farihi09,xu12,hoard13}). The discovery of these four new debris disks enriches the current dusty WD population with one of the coolest (J1507), hottest (J1537), and the most massive (J1234) WDs known to host circumstellar dust. Expanding the parameter space known to be hospitable to circumstellar dust at WD stars will not only guide future infrared excess searches, but also enhance our understanding of the formation and evolution of these remnant planetary systems."
    },
    "1402/1402.6385_arXiv.txt": {
        "abstract": "\\noindent We present the results of a time series analysis of the long-term radio lightcurves of four blazars: 3C 279, 3C 345, 3C 446, and BL Lacertae. We exploit the data base of the University of Michigan Radio Astronomy Observatory (UMRAO) monitoring program which provides densely sampled lightcurves spanning 32 years in time in three frequency bands located at 4.8, 8, and 14.5\\,GHz. Our sources show mostly flat or inverted (spectral indices $-0.5\\lesssim\\alpha\\lesssim0$) spectra, in agreement with optically thick emission. All lightcurves show strong variability on all time scales. Analyzing the time lags between the lightcurves from different frequency bands, we find that we can distinguish high-peaking flares and low-peaking flares in accord with the classification of \\cite{Valtaoja}. The periodograms (temporal power spectra) of the observed lightcurves are consistent with random-walk powerlaw noise without any indication of (quasi-)periodic variability. The fact that all four sources studied are in agreement with being random-walk noise emitters at radio wavelengths suggests that such behavior is a general property of blazars. ",
        "introduction": "}  \\noindent The strong and complex temporal flux variability of Active Galactic Nuclei (AGN; see, e.g., \\citealt{Beckmann} and references therein for a recent review) provides valuable information on the internal conditions of accretion zones and plasma outflows. Various characteristic variability patterns have been associated with a wide range of physical phenomena, from shocks in continuous (e.g., \\citealt{Marscher}) or discontinuous (e.g., \\citealt{spada2001}) jets to orbiting plasma ``hotspots'' (e.g., \\citealt{abramowicz1991}) or plasma density waves (e.g., \\citealt{kato2000}) in accretion disks. Accordingly, multiple studies have aimed at quantifying the properties of AGN variability on all time scales and throughout the electromagnetic spectrum. At radio wavelengths, variability time scales probed by observations range from tens of minutes (\\citealt{Schodel}, studying the mm/radio lightcurve of M 81*; see also \\citealt{Kim2013} for a discussion of the detectability of intra-day variability) to tens of years (\\citealt{Hovatta2007}, in a statistical analysis of the long-term flux variability of 80 AGN). Of particular interest is the possible presence of quasi-periodic oscillations (QPO) which has been reported by several studies of blazar lightcurves (e.g. \\citealt{Rani2009, Rani2010, Gupta2012}).  Fourier transform, period folding, power spectrum, and periodogram methods (cf. \\citealt{priestley1981} for an exhaustive review of time series analysis) have been used extensively for quantifying the statistical properties of AGN variability and for the search for possible QPOs (e.g. \\citealt{Benlloch} for X-ray, \\citealt{Webb} for optical, \\citealt{Fan1999} for near infrared, and \\citealt{Aller} for radio observations). As already noted by \\citet{Press}, power spectra of AGN lightcurves globally follow power laws $A_f{\\propto}f^{-\\beta}$ with $\\beta>0$, corresponding to \\emph{red noise};\\footnote{In the context of time series analysis, the term ``noise'' refers to stochastic emission from a source of radiation, \\emph{not} to measurement errors or instrumental noise.} here $A_f$ denotes the power spectral amplitude as function of sampling frequency $f$. Lightcurves composed of pure Gaussian \\emph{white noise} have flat power spectra ($\\beta=0$). Other important special cases are \\emph{random walk noise} ($\\beta = 2$) -- which is the integral of white noise -- and \\emph{flicker noise} ($\\beta = 1$, ``$1/f$ noise'') as intermediate case between white noise and random walk noise (see also \\citealt{Park2012} for a detailed technical discussion). \\citet{Lawrence1987} found that the power spectrum of the Seyfert galaxy NGC 4051 can be described by flicker noise. Red noise power spectra were also observed by \\cite{Lawrence1993} who used 12 high-quality ``long look'' X-ray lightcurves of AGN.   \\begin{deluxetable*}{ccccccccc}  \\tablecaption{Properties of our four target blazars \\label{Information}} \\tablehead{ \\colhead{Object} & \\colhead{RA} & \\colhead{DEC} & \\colhead{Type} & \\colhead{Redshift} & \\colhead{$T$ [yr]} & \\colhead{$N_{4.8}$} & \\colhead{$N_{8.0}$} & \\colhead{$N_{14.5}$} } \\startdata 3C 279 & 12:56:11 & $-$05:47:22 & FSRQ   & 0.536 & 32.52 & 1086 & 1337 & 1473 \\\\ 3C 345 & 16:42:59 & +39:48:37   & FSRQ   & 0.593 & 32.54 & 1323 & 1315 & 1415 \\\\ 3C 446 & 22:25:47 & $-$04:57:01 & BL Lac & 1.404 & 32.12 & 680  & 902  & 1088 \\\\ BL Lac & 22:02:43 & +42:16:40   & BL Lac & 0.069 & 32.54 & 1256 & 1315 & 1755 \\enddata  \\tablecomments{J2000 coordinates, source types, and redshifts are taken from the NED. We also give the total monitoring time $T$ (in years) and the numbers $N$ of flux data points for 4.8, 8.0, and 14.5\\,GHz, respectively.}  \\end{deluxetable*}   A multitude of studies illustrates the difficulties of determining the statistical significance of supposed QPO signals in the power spectra of AGN lightcurves. The analysis of \\cite{Benlloch} concluded that a previously reported quasi-periodic signal in X-ray lightcurves of the Seyfert galaxy Mrk 766 was actually statistically insignificant. \\cite{Uttley} pointed out the importance of sampling effects leading to \\emph{red-noise leaks} and \\emph{aliasing}. \\cite{Vaughan} gives an analytical approach to derive significance levels for peaks in red-noise power spectra. \\cite{Do} demonstrated the power of Monte-Carlo techniques for deriving significance levels by comparing the red-noise power spectra of actual and simulated flux data.  The temporal flux variability of AGN can be exploited for elucidating the physical conditions within active galaxies especially at radio frequencies where monitoring observations of hundreds of targets have been conducted over several decades by various observatories. Remarkably, many studies aimed at analyzing long-term AGN radio variability do not take into account the intrinsic red-noise properties of the lightcurves. The incorrect assumption of constant (as function of $f$) significance levels in power spectra (following from the assumption of white-noise dominated lightcurves) has lead to reports of ``characteristic'' time scales which are actually not special at all (cf., e.g., \\citealt{Ciaramella, Hovatta2007, Nieppola}).  Blazars, characterized by violent flux variability across the entire electromagnetic spectrum, are a subset of AGN which include BL Lacertae (BL Lac) objects and Flat Spectrum Radio Quasars (FSRQs). In accordance with the standard viewing angle unification scheme of AGN \\citep{Urry}, it is commonly assumed that their observed emission is generated by synchrotron radiation -- dominating from radio to optical frequencies -- and inverse Compton emission -- dominating at frequencies higher than optical -- from relativistic plasma jets (almost) aligned with the line of sight. In order to perform a thorough study of the statistical properties of blazar emission, we analyze the lightcurves of four radio-bright blazars with strong flux variability -- 3C 279, 3C 345, 3C 446, and BL Lac -- provided by the University of Michigan Radio Astronomy Observatory (UMRAO) monitoring program of AGN. The data set comprises data spanning $\\approx$32 years in time and covering three frequency bands located at 4.8\\,GHz, 8.0\\,GHz, and 14.5\\,GHz. ",
        "conclusions": "  \\begin{enumerate}  \\item  Our sources show mostly flat or inverted ($-0.5\\lesssim\\alpha\\lesssim0$) spectral indices, in agreement with optically thick synchrotron emission. The lightcurves of different frequencies are either simultaneous (within errors) or shifted relative to each other such that the high-frequency emission leads the low-frequency emission by up to $\\approx$1.5 years. We are able to distinguish high-peaking and low-peaking flares according to the classification of \\cite{Valtaoja}.  \\item  All lightcurves show variability on all time scales. Their periodograms (power spectra) are in agreement with being pure red-noise powerlaw spectra without any indication for (quasi-)periodic signals. When taking into account the sampling patterns via dedicated Monte Carlo simulations, we find that all lightcurves are consistent with being random walk noise signals with powerlaw slopes $\\beta\\approx2$. Given that we find this behavior in all four sources under study, this suggests that random walk noise lightcurves are a general feature of blazars.  \\end{enumerate}  Our results imply that careful time series analysis of high-quality blazar lightcurves provides information on the source structure even if a target is not resolved spatially. Obviously, it will be necessary to systematically study much larger blazar samples in order to decide if the trends we have uncovered are indeed general."
    },
    "1402/1402.5948_arXiv.txt": {
        "abstract": "Fourier analysis of the light curve of AC And from the HATNet database reveals the rich frequency structure of this object. Above $30$ components are found down to the amplitude of $3$~mmag. Several of these frequencies are not the linear combinations of the three basic components. We detect period increase in all three components that may lend support to the Pop I classification of this variable. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.2103_arXiv.txt": {
        "abstract": "We present new observations of the XZ~Tau system made at high angular resolution (55\\,milliarcsec) with the Karl G.\\ Jansky Very Large Array (VLA) at a wavelength of 7\\,mm.  Observations of XZ~Tau made with the VLA in 2004 appeared to show a triple-star system, with XZ~Tau~A resolved into two sources, XZ~Tau~A and XZ~Tau~C.  The angular separation of XZ~Tau~A and C (0.09\\,arcsec) suggested a projected orbital separation of around 13\\,AU with a possible orbital period of around 40\\,yr.  Our follow-up observations were obtained approximately 8\\,yr later, a fifth of this putative orbital period, and should therefore allow us to constrain the orbital parameters of XZ~Tau~C, and evaluate the possibility that a recent periastron passage of C coincided with the launch of extended optical outflows from XZ~Tau~A. Despite improved sensitivity and resolution, as compared with the 2004 observations, we find no evidence of XZ~Tau~C in our data.  Components A and B are detected with a signal-to-noise ratio greater than ten; their orbital motions are consistent with previous studies of the system, although the emission from XZ~Tau~A appears to be weaker.  Three possible interpretations are offered: either XZ~Tau~C is transiting XZ~Tau~A, which is broadly consistent with the periastron passage hypothesis, or the emission seen in 2004 was that of a transient, or XZ~Tau~C does not exist.  A fourth interpretation, that XZ~Tau~C was ejected from the system, is dismissed due to the lack of angular momentum redistribution in the orbits of XZ~Tau~A and XZ~Tau~B that would result from such an event.  Transients are rare but cannot be ruled out in a T~Tauri system known to exhibit variable behaviour.  Our observations are insufficient to distinguish between the remaining possibilities, at least not until we obtain further VLA observations at a sufficiently later time.  A further non-detection would allow us to reject the transit hypothesis, and the periastron passage of XZ~Tau~C as agent of XZ Tau A's outflows. ",
        "introduction": "XZ~Tau is a binary system composed of a T~Tauri star, XZ~Tau~A, with a cool companion, XZ~Tau~B, separated by approximately 0.3\\,arcsec, at a distance of approximately 140\\,pc from Earth \\citep{Haas1990,Kenyon1994,Torres2009}.  Like many other T~Tauri stars, XZ~Tau~A drives collimated jets and optical outflows \\citep{Mundt1990,Krist1997}.  {\\it Hubble Space Telescope} imaging of these outflows shows nebular emission in the shape of an elongated bubble with expansion velocities of around 70\\,$\\mathrm{km\\,s^{-1}}$ \\citep{Krist1999}.  The substructure displayed by the bubble suggests its driver is episodic, with the cause attributed to a velocity pulse in the jet of XZ~Tau~A, triggered in the early 1980s \\citep{Krist2008}.  These previous studies, particularly that of \\citeauthor{Krist2008}, explored the possibility that the periastron passage of XZ~Tau~B could have caused the outflows (cf.\\ \\citealt{enc_outburst}).  However, this would require an eccentric orbit, which is inconsistent with observations of the A/B system, which instead suggest a circular, face-on orbit.  Also, the periastron passage of XZ~Tau~B would have occurred in the 1950s, too early to cause the outflow.  The periastron passage hypothesis was revived by more recent observations of the XZ~Tau system using the Very Large Array (VLA) by \\citet{Carrasco-Gonzalez2009}.  Observations of the 7-mm continuum resolved XZ~Tau~A into two components, with the new component, XZ~Tau~C, separated by around 0.09\\,arcsec (13\\,AU).  The non-detection of XZ~Tau~C in the optical waveband was suggestive of a stellar object heavily embedded in a dusty envelope or disk.  While a single detection yielded no information on the orbit of component C around A, the existence of XZ~Tau~C increased the likelihood of a close approach to XZ~Tau~A, making it a potential trigger for the outflow.  \\citet{Carrasco-Gonzalez2009} speculate that if the orbit of XZ~Tau~C is circular, and the total system mass is $\\approx$1\\,M$_\\odot$, then the orbital period of XZ~Tau~C should be around 40\\,yr. If the A/C system was close to apastron at the epoch of detection (2004), this would place XZ~Tau~C at periastron in the 1980s, as required, and a further ejection from XZ~Tau~A could be expected around 2020.  However, the data available to the authors was insufficient to confirm orbital parameters, and as such this explanation for the outflows was tentative.  Further observations of XZ~Tau~C at sufficiently later epochs would be required to either confirm or refute the periastron passage model for outflow generation.  To this end, we observed the XZ~Tau system at high angular resolution, using the newly upgraded Karl G.\\ Jansky VLA at 7\\,mm, to confirm the existence of XZ~Tau~C and constrain its orbital parameters.  The time interval between our observations (2012) and the previous observations (2004) corresponds to approximately one fifth of the potential orbital period of XZ~Tau~C.  These observations yield no detection of XZ~Tau~C, and the positions of XZ~Tau~A and B are consistent with the orbital solutions presented by previous studies.  This Letter is composed as follows: we describe the observations taken in \\S\\ref{sec:observations}; we discuss the results in \\S\\ref{sec:results}, and we summarise the work in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  \\noindent We report observations of the XZ~Tau system, recently observed to possess a third component, XZ~Tau~C \\citep{Carrasco-Gonzalez2009}.  This third component had been proposed as a potential driver for the outflows generated by XZ~Tau~A, which could have been generated by a periastron passage of this new object in the 1980s.  Our new observations were made using the VLA in its most extended configuration, after a period of roughly one fifth of the object's assumed orbital period, and can thereby constrain the orbit of XZ~Tau~C and test the periastron passage theory.  However, our observations yield no detection of XZ~Tau~C.  XZ~Tau~A and B are both detected with a high degree of confidence, with positions and orbits consistent with those described in the literature.  Three potential interpretations of the data are possible.  Of these, the most prosaic is that XZ~Tau~C does not exist, and was an artefact of reducing interferometric data, which is consistent with its non-detection in the optical. Alternatively, XZ~Tau~C may have been caught whilst transiting XZ~Tau~A, which is consistent with the orbital requirements for XZ~Tau~C to trigger outflows at periastron passage.  This second interpretation would require the hierarchical triple to possess a large mutual inclination, and hence be dynamically unstable, but on timescales that are sufficiently long to remain consistent with the observations. Finally, we cannot rule out the possibility of a transient event local to this T~Tauri system, a system known to be variable. To distinguish between these potential interpretations, determining the nature of XZ~Tau~C, requires a very brief, future observation using the VLA in A~configuration.  If the transit hypothesis is true, future observations should be able to detect XZ~Tau~C once it has moved away from XZ~Tau~A.  If observations at this time fail to detect XZ~Tau~C, this would be sufficient to reject it as a cause of outflows being generated by XZ~Tau~A."
    },
    "1402/1402.2784_arXiv.txt": {
        "abstract": "{ We investigate the Euclidean vacuum mode functions of a massive vector field in a spatially open chart of de Sitter spacetime. In the one-bubble open inflationary scenario that naturally predicts a negative spatial curvature after a quantum tunneling, it is known that a light scalar field has the so-called supercurvature mode, i.e. an additional discrete mode which describes fluctuations over scales larger than the spatial curvature scale. If such supercurvature modes exist for a vector field with a sufficiently light mass, then they would decay slower and easily survive the inflationary era. However, the existence of supercurvature mode strongly depends on details of the system. To clarify whether a massive vector field has supercurvature modes, we consider a U(1) gauge field with gauge and conformal invariances spontaneously broken through the Higgs mechanism, and present explicit expressions for the Euclidean vacuum mode functions. We find that, for any values of the vector field mass, there is no supercurvature mode. In the massless limit, the absence of supercurvature modes in the scalar sector stems from the gauge symmetry. }     \\begin{document} ",
        "introduction": "Recent observational data provide a strong support of the existence of extragalactic magnetic fields, in the range of $\\mcO (10^{-14}$-$10^{-20})\\,{\\rm G}$ on Mpc scales~\\cite{Neronov:1900zz,Tavecchio:2010mk,Dermer:2010mm,Huan:2011kp,Dolag:2010ni,Essey:2010nd,Taylor:2011bn,Vovk:2011aa,Takahashi:2011ac,Finke:2013bua}\\,. The generation of the magnetic field in high-redshift galaxies, clusters, and even in empty intergalactic region is still an unresolved problem in cosmology. No promising astrophysical process to generate the sufficient amount of the magnetic field on the large scales are known. As for the inflationary magnetogenesis, though the various mechanism are proposed, the several difficulties such as the strong coupling problem, the backreaction problem and the curvature perturbation problem in some specific models prevent successful production of magnetic field~\\cite{Demozzi:2009fu, Barnaby:2012xt, Fujita:2013qxa}. Actually both upper and lower limits on the inflation energy scale can be derived from these problems in model independent ways and the limits are considerably severe if the extragalactic magnetic fields are stronger than $10^{-16}$G at present~\\cite{Fujita:2012rb,Suyama:2012wh,Fujita:2014sna}. Thus it is known to be very difficult to generate the magnetic field in the context of the inflationary magnetogenesis on the flat Friedmann-Lema\\^itre-Robertson-Walker (FLRW) universe.   The superadiabatic growth of the magnetic fields in the open FLRW universe has been discussed in the literatures~\\cite{Barrow:2011ic,Barrow:2012ty}. The authors of these literatures assumed the existence of supercurvature modes of the magnetic field, which describes the fluctuations with the wavelength exceeding the spatial curvature scale. If a supercurvature mode exists, it decays slower than $1/a^2$\\,, where $a$ corresponds to the conventional scale factor of a FLRW universe, and can easily survive the inflationary era. Hence the relatively large amount of the magnetic field on supercurvature scales would remain at late time. However, the existence of supercurvature modes of magnetic fields is non-trivial and should be critically studied. Adamek {\\it et al.}~\\cite{Adamek:2011hi} recently pointed out that the equations of motion of a U(1) gauge field with unbroken conformal and gauge symmetries can be rewritten in the form that is identical to those of massive scalar fields for which there is no supercurvature mode~\\footnote{Rigorously speaking, the proof of the absence of supercurvature modes requires knowledge of not only the equation of motion but also a Klein-Gordon norm and proper boundary conditions. The present paper fills those gaps for the analysis of the massless vector field, although our main focus will be on a massive vector field.}.   The purpose of the present paper is to investigate whether supercurvature modes exist for a massive vector field, in both scalar and vector sectors of the physical spectrum. To be specific, we consider a U(1) gauge field with both gauge and conformal symmetries spontaneously broken through the Higgs mechanism. As for the background geometry, we consider a de Sitter spacetime in the open chart. This is relevant to the one-bubble open inflation scenario that naturally predicts the spatially negative curvature universe. While the recent observational data show that the universe is almost exactly flat with accuracy of about $1\\%$\\,, $|1-\\Omega_0 |\\leq 10^{-2}$~\\cite{Ade:2013zuv}\\,, open inflation scenario is attracting a renewed interest in the context of the string landscape scenario~\\cite{Susskind:2003kw,Freivogel:2004rd}. There are a huge number of metastable de Sitter vacua and the tunneling transition generally occurs through the nucleation of a true vacuum bubble in the false vacuum background. Because of the symmetry of the instanton solution, a bubble formed by the Coleman-De Luccia (CDL) instanton~\\cite{Coleman:1977py,Coleman:1980aw} looks like an infinite open universe from the viewpoint of an observer inside. If the universe experienced a sufficiently long inflation after the bubble nucleation, then the universe becomes almost exactly flat and subsequently evolves as a slightly open FLRW universe. This leads to a natural realization of one-bubble open inflation (see e.g. \\cite{Sasaki:1994yt,Yamamoto:1996qq,Tanaka:1997kq,Garriga:1998he,Garriga:1997wz}) and can be tested against observations~\\cite{Yamauchi:2011qq,Sugimura:2012kr,Sugimura:2013cra}.    This paper is organized as follows. We first illustrate the background spacetime in section \\ref{sec:Background}. In section \\ref{sec: KG norm}\\,, we expand the U(1) gauge field by harmonic functions and write down the reduced action for the even and odd modes of the U(1) gauge field. In order to investigate the existence/absence of supercurvature modes, we show the quantization conditions for the even and odd modes on a Cauchy surface. With the obtained normalization conditions, we then analyze whether the supercurvature modes, which are normalizable on the Cauchy surface, exist in section \\ref{sec: Mode functions}.  In section \\ref{sec: massless limit}\\,, as a consistency check, we explicitly calculate the Wightman function in the decoupling limit by using the (subcurvature) mode functions derived in section \\ref{sec: KG norm}. It is shown that the correct expression for the Euclidean Wightman function is recovered in the decoupling limit without need for any supercurvature modes. Finally, section \\ref{sec:summary} is devoted to a summary and discussions. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have investigated the Euclidean vacuum mode functions of a massive vector field in the spatially open chart of de Sitter spacetime. In order to clarify whether supercurvature modes exist, we have studied the U(1) gauge field with gauge and conformal symmetries spontaneously broken through the Higgs mechanism. We have found that there is no supercurvature mode for both the even and odd parity sectors. This implies that it is difficult to generate the sufficient amount of  the magnetic field on large scales by using the superadiabatic growth within the one-bubble open inflation scenario even if the Higgs mechanism spontaneously breaks gauge and conformal invariances.   Utilizing the obtained mode functions, we have explicitly computed the Wightman function of the U(1) gauge field in terms of the coordinates in the open chart of the de Sitter spacetime, and have compared it with one obtained by other methods. It was found that the leading-order Wightman function in the decoupling ($e\\to 0$) limit is correctly reproduced by the sum of the products of the subcurvature modes without need for introducing supercurvature modes. In consequence we have verified that the supercurvature mode is not needed as a part of a complete set of mode functions of the U(1) gauge field in the decoupling limit.   An interesting observation made in subsection \\ref{sec:exact massless} is that the existence/absence of supercurvature modes can be strongly related to symmetries of the theory. While a massive scalar field with a sufficiently light mass has a supercurvature mode~\\cite{Sasaki:1994yt} that survives the massless limit, a theory of a scalar field with shift symmetry does not allow a physical supercurvature mode. This is because the would-be supercurvature mode does not show up in any correlation functions invariant under the shift symmetry. Furthermore, a vector field with a U(1) gauge symmetry does not have a supercurvature mode even when the vector field is given a mass by the Higgs mechanism and thus absorbs a light scalar degree of freedom (the phase of the complex Higgs field). It may be interesting to investigate supercurvature modes of the vector field when we take metric perturbations into account since gravity has the diffeomorphism symmetry, although in the present paper we take account of the effect of gravity only through a curved background. The evaluation of the metric perturbations is beyond the scope of the present paper and we hope to come back to this issue in a future publication.   In this paper, we have assumed several simplifications: (i) the universe during inflationary era after a quantum tunneling is assumed to be well approximated by an exact de Sitter spacetime in the open chart; (ii) the origin of the breaking of the gauge and conformal symmetries and the mass of the vector field is assumed to be the standard Higgs mechanism; (iii) the mass squared $V''_{\\Phi}(|\\Phi|)$ around the minimum of the Higgs potential is assumed to be large enough so that the mass of the vector field can be considered as constant during inflation. Relaxing some of these assumptions would in principle affect details of our results, although generic features that we have found are expected to remain the same. Furthermore, we have neglected the interactions between the tunneling field and the other fields such as the Higgs field. If we take into account such interactions, then spatially localized, bubble-shaped features may appear~\\cite{Sugimura:2012kr}. We hope to come back to these issues in the near future."
    },
    "1402/1402.7168_arXiv.txt": {
        "abstract": "{The results of a search for eclipsing Am star binaries using photometry from the SuperWASP survey are presented. The light curves of 1742 Am stars fainter than $V$ = 8.0 were analysed for the presences of eclipses. A total of 70 stars were found to exhibit eclipses, with 66 having sufficient observations to enable orbital periods to be determined and 28 of which are newly identified eclipsing systems. Also presented are spectroscopic orbits for 5 of the systems. The number of systems and the period distribution is found to be consistent with that identified in previous radial velocity surveys of `classical' Am stars.} ",
        "introduction": "\\label{introduction}  Amongst the A and F stars there exists a subclass of peculiar stars called the metallic-lined (Am) stars, in which the \\ion{Ca}{II}~K line is considerably weaker than would be expected from the average metallic line type \\citep{1940ApJ....92..256T,1948ApJ...107..107R}. These stars exhibit an apparent underabundance of calcium and scandium, overabundances of iron-group elements, and extreme enhancements of rare-earth elements \\citep{1970PASP...82..781C}. In contrast to normal A-type stars, the Am stars are slowly rotating \\citep{1983aspp.book.....W} with maximum $v \\sin i$ values of $\\sim$100\\,km\\,s$^{-1}$ \\citep{1973ApJ...182..809A}. The abundance anomalies are thought to be due to radiative diffusion of elements within the stable atmospheres of these relatively slowly rotating stars \\citep{1970ApJ...160..641M,1980AJ.....85..589M,1983ApJ...269..239M}.  Early spectral studies of Am stars hinted at a high fraction of spectroscopic binaries \\citep{1948ApJ...107..107R}, while the systematic study by \\cite{1961ApJS....6...37A} led to the conclusion that all Am stars are members of spectroscopic binaries. Hence, it was assumed that the slow rotation of Am stars, required for radiative diffusion to occur, was the result of the reduction of rotational velocities due to tidal interaction. While there were spectroscopic orbits for many Am stars \\citep[e.g.][]{2004A&A...424..727P}, only a handful were known to be eclipsing \\citep[e.g.][]{1980ARA&A..18..115P,1991A&ARv...3...91A}. It was this that led \\cite{1990clst.book.....J} to conclude that: \\begin{quote} \\emph{A curious fact is that among the many Am stars known (all of which are binaries) there should be many eclipsing binaries, but surprisingly very few cases are known.} \\end{quote}  Comprehensive spectroscopic radial velocity studies of carefully selected Am stars, have found a binary fraction of nearer 60--70\\% \\citep{1985ApJS...59..229A,2007MNRAS.380.1064C}. The period distribution shows that the majority of systems have periods $\\la$\\,50\\,days, consistent with the slow rotation being due to tidal synchronisation or pseudo-synchronization \\citep{1996A&A...313..523B,1997A&A...326..655B}. There are, nonetheless, systems with longer periods, suggesting that these Am stars were formed with low initial rotation velocities \\citep{2004ASPC..318..297N}. The key to the Am phenomenon appears to be slow rotation and not binarity per se.  The recent \\citet{2009A&A...498..961R} catalogue lists only 61 Am stars as eclipsing or possibly eclipsing. This represents only 1.4\\% of the Am stars in the catalogue. Given that a large fraction of Am stars are supposed to be in binary systems, this percentage does appear rather low. For example, in a binary with a period of $\\sim$5\\,days (typical of many Am spectroscopic binaries), there is a $\\sim$10\\% probability that the system should be eclipsing. Hence, there is a perhaps somewhat na\\\"{i}ve expectation that there ought to be more eclipsing Am stars. It is this which led us to investigate the number of eclipsing systems that can be found using light curves obtained from the SuperWASP exoplanet transit survey. ",
        "conclusions": "\\subsection{Expected period distribution of Am binaries}  The results from the radial velocity studies of \\cite{1985ApJS...59..229A} and \\cite{2007MNRAS.380.1064C} can be used to predict the number and period distribution of eclipsing Am binaries. The combined sample comprises 151 Am stars, with 61 SB1s and 28 SB2s. This was binned onto 1-day bins, normalized by the total number of stars, to generate a period probability distribution for Am stars. Multiplying by the estimated WASP detection probability ($p_{\\rm overall}$) obtained in Sect.\\,\\ref{DetProb} and by the number of stars in the WASP sample (1742), yields an estimate of the expected period distribution of eclipsing Am stars. In Figure\\,\\ref{dist-predicted} the WASP eclipsing Am star period distribution is compared to that predicted for the two identical stars and the dark companion cases. Since these represent the extrema of the probabilities, we also include the predicted distribution obtained using the ratio of SB1 and SB2 systems from the RV studies.  The number of eclipsing Am stars found by WASP does appear to be broadly consistent with the expected number of systems. We recall from Sect.\\,\\ref{introduction} that the fraction of spectroscopic binaries is 60$\\sim$70\\%. Thus, the eclipsing fraction appears to be similar, suggesting a significant fraction of Am stars might be single or have hard to detect companions. The period distribution is, however, slightly different, with a pronounced peak at shorter periods due to the inclusion of close binaries. The $v \\sin i$ distributions of both \\cite{1985ApJS...59..229A} and \\cite{2007MNRAS.380.1064C} are skewed toward lower values than the \\citet{2009A&A...498..961R} sample. These radial velocity studies have preferentially avoided stars with high rotation, which accounts for the excess of short period systems found in the WASP sample. The distribution of Am-type spectroscopic binaries in the \\citet{2009A&A...498..961R} catalogue (some 210 systems) shows a similar short orbital period excess, due to the inclusion of Am stars with a wide range of rotational velocities.  \\begin{figure}[h] \\includegraphics[width=\\columnwidth]{dist-predicted.eps} \\caption{Period distribution of eclipsing Am star binaries. The WASP eclipsing Am star distribution is given as solid grey. The predicted period distribution based on the results of spectroscopic binaries are given as dashed-line for the two identical stars and dotted-line for the dark companion case. The thicker solid-line is that predicted based on the ratio of SB1 and SB2 systems.} \\label{dist-predicted} \\end{figure}  \\subsection{Mass ratio distribution}  \\begin{table*} \\caption{Results from \\textsc{jktebop} fits to light curves (Detached systems only) and approximate stellar parameters, assuming primary is ${T_\\mathrm{eff}}_A = 7500$\\,K. See Text for details.} \\label{JKTEBOP} \\begin{tabular}{llllllllllll}\\hline\\hline Renson &$\\frac{R_A+R_B}{a}$ & $R_B/R_A$ & $i$ & $e$ & $J_B/J_A$ & $P$ & ${T_\\mathrm{eff}}_B$ & $R_A$ & $R_B$ & $q$ & $q$ (literature) \\\\ \\hline 3750& 0.307& 0.111& 89.95&  0.00&  0.00&  1.2199& 2000\\S&  1.63&  0.18&  0.07 &   $<$0.003    {\\citep{2010MNRAS.407..507C}}\\\\ 4660& 0.298& 0.910& 87.61&  0.00&  0.84&  2.7807& 7176&  1.97&  1.80&  0.92 &   0.91        {\\citep{2011MNRAS.414.3740S}}\\\\ 6720& 0.344& 1.354& 79.59&  0.00&  0.91&  2.3832& 7333&  1.68&  2.26&  1.07 &\\\\ 7310& 0.161& 0.601& 86.37&  0.00&  0.50&  6.6637& 6322&  2.20&  1.32&  0.71 &\\\\ 7730& 0.156& 3.000& 89.57&  0.24&  0.10& 39.2827& 4218&  2.93&  8.71&  0.84 &   0.96        {\\citep{1988AJ.....96.1040P}}\\\\ 8215& 0.105& 0.940& 86.43&  0.00&  0.44&  8.7959& 6126&  1.38&  1.29&  0.77&\\\\ 9237& 0.248& 1.036& 80.52&  0.00&  0.49&  2.8562& 6288&  1.51&  1.55&  0.81&\\\\ 9318& 0.148& 0.472& 85.89&  0.06&  0.72& 11.1133& 6901&  3.24&  1.52&  0.71&\\\\ 9410& 0.189& 0.634& 87.74&  0.02&  0.30& 16.7873& 5537&  5.07&  3.19&  0.59&   0.81{\\dag}  {\\citep{1971AJ.....76..544L}}\\\\ 10016& 0.242& 1.323& 86.98&  0.25&  1.03&  5.4309& 7557&  2.11&  2.81&  1.11&\\\\ 11100& 0.165& 0.397& 85.07&  0.00&  0.36&  4.0375& 5818&  1.79&  0.71&  0.61&\\\\ 11470& 0.228& 0.580& 89.36&  0.00&  0.39&  2.8721& 5931&  1.75&  1.02&  0.66&\\\\ 14850& 0.215& 0.272& 83.02&  0.00&  0.08&  3.9795& 3925&  2.43&  0.66&  0.30&\\\\ 15034& 0.202& 0.469& 83.88&  0.00&  0.54&  7.5392& 6423&  3.39&  1.59&  0.65&\\\\ 15190& 0.174& 0.439& 82.60&  0.00&  0.02&  5.1229& 2952&  1.97&  0.86&  0.21&   0.33{\\dag}  {\\citep{2003MNRAS.346..555C}}\\\\ 15445& 0.270& 0.881& 86.32&  0.00&  0.92&  5.7606& 7355&  3.10&  2.72&  0.93&\\\\ 21400& 0.106& 0.513& 86.63&  0.00&  0.04&  7.7729& 3435&  1.49&  0.77&  0.29&\\\\ 22860& 0.136& 0.558& 83.84&  0.28&  2.81&  7.9006& 9712&  2.33&  1.30&  1.17&\\\\ 25020& 0.188& 0.081& 89.10&  0.00&  0.00&  6.6284& 2000\\S&  3.39&  0.28&  0.06&\\\\ 25070& 0.185& 0.500& 85.60&  0.14&  0.91&  6.7149& 7323&  2.85&  1.42&  0.79&   1.00 (this work)\\\\ 29290& 0.314& 0.977& 79.45&  0.00&  0.95&  3.2922& 7406&  2.30&  2.23&  0.98&\\\\ 30090& 0.221& 0.788& 82.26&  0.00&  0.97&  4.3486& 7451&  2.11&  1.67&  0.93&\\\\ 30110& 0.148& 0.162& 84.55&  0.00&  0.09&  4.3145& 4108&  1.92&  0.31&  0.37&\\\\ 30457& 0.137& 0.231& 86.82&  0.00&  0.00& 11.9420& 2000\\S&  3.20&  0.74&  0.08&\\\\ 30650& 0.060& 0.378& 88.38&  0.00&  0.00& 18.1210& 2000\\S&  1.53&  0.58&  0.09&\\\\ 35000& 0.190& 0.835& 82.07&  0.00&  0.00& 10.2861& 2000\\S&  2.69&  2.26&  0.13&\\\\ 36660& 0.096& 0.698& 88.61&  0.41&  0.71& 16.3653& 6892&  2.30&  1.60&  0.81&   0.80 (this work)\\\\ 37220& 0.135& 0.962& 85.57&  0.00&  0.97&  5.7931& 7437&  1.38&  1.33&  0.98&\\\\ 37610& 0.240& 1.455& 85.50&  0.00&  1.28&  3.2387& 7974&  1.38&  2.02&  1.19&\\\\ 38500& 0.173& 0.348& 83.65&  0.00&  0.19&  3.9938& 4945&  1.87&  0.65&  0.47&\\\\ 40350& 0.156& 0.272& 88.22&  0.00&  0.00&  7.0687& 2000\\S&  2.43&  0.66&  0.08&\\\\ 40910& 0.132& 1.114& 87.17&  0.05&  1.05& 11.1163& 7594&  2.03&  2.23&  1.05&\\\\ 42906& 0.062& 0.860& 87.84&  0.00&  0.00& 19.6987& 2000\\S&  1.23&  1.06&  0.12&   0.98        {\\citep{2003MNRAS.346..555C}}\\\\ 44140& 0.309& 1.391& 86.51&  0.02&  0.66&  2.0598& 6763&  1.31&  1.81&  0.96&   0.90        {\\citep{1970ApJ...162..925P}}\\\\ 49380& 0.161& 0.355& 85.06&  0.11&  0.57&  5.0204& 6508&  2.14&  0.76&  0.68&   0.83 (this work)\\\\ 51506& 0.322& 0.119& 80.30&  0.12&  0.17&  1.9196& 4787&  2.69&  0.32&  0.46&   0.12 (this work)\\\\ 56310& 0.285& 1.108& 84.26&  0.00&  0.88&  2.6959& 7258&  1.68&  1.85&  0.99&\\\\ 58170& 0.366& 0.854& 88.24&  0.00&  0.48&  1.6047& 6246&  1.65&  1.42&  0.77&   0.78        {\\citep{1968ApJ...154..191P}}\\\\ 58256& 0.271& 0.378& 81.63&  0.00&  0.41&  2.9673& 5983&  2.50&  0.95&  0.59&\\\\ 59780& 0.228& 1.200& 87.87&  0.33&  0.95&  7.3514& 7412&  2.61&  3.15&  1.06&   0.98        {\\citep{1999AJ....118.1831T}}\\\\ \\hline \\end{tabular} \\tablefoot{\\dag Mass ratio (q) obtained using spectroscopic binary mass function, $f(m)$, and assuming $M_1$ = 1.7\\,$M_\\sun$ and $i$ = 90.\\newline \\S lower-limit on {${T_\\mathrm{eff}}_{B}$} imposed when $J_B/J_A = 0$. } \\end{table*}  Without direct determinations of masses from spectroscopic studies, we can only make a rather crude estimate of the mass distribution of the eclipsing systems from their light curves and the \\textsc{jktebop} fits. Since the bolometric correction for late-A stars is small, we can make the approximation that the ratio of bolometric surface brightnesses is given by the WASP bandpass surface brightness ratio ($J_B/J_A$). Thus the effective temperature of the secondary (${T_\\mathrm{eff}}_{B}$) can be obtained from, \\[{T_\\mathrm{eff}}_{B} \\approx {T_\\mathrm{eff}}_{A} \\times (J_B/J_A)^{1/4},\\] where the effective temperature of the primary (${T_\\mathrm{eff}}_{A}$) is assumed to be 7500\\,K. With initial mass estimates of $M_A$ = $M_B$ = 1.7, the known orbital period ($P_{\\rm orb}$) and sum of the radii ($\\frac{R_A+R_B}{a}$) from \\textsc{jktebop}, we determine initial values for $R_A + R_B$. Using the ratio of the radii also from \\textsc{jktebop} ($R_B/R_A$) we can determine individual values for $R_A$ and $R_B$. Next, the \\cite{2010A&ARv..18...67T} relations are used to determine stellar mass estimates, by varying $\\log g$ so that radii obtained from \\cite{2010A&ARv..18...67T} agrees with that expected from the \\textsc{jktebop} analysis. The procedure is iterated until there are no changes in parameters. The results are presented in Table\\,\\ref{JKTEBOP}, along with values for mass ratio ($q$) from spectroscopic analyses in the literature and those determined in Section\\,\\ref{Data_modelling}. The average difference between our estimated $q$ values and the spectroscopic values is $-$0.11, but with an rms scatter of $\\pm$0.25.  \\begin{figure}[h] \\includegraphics[width=\\columnwidth]{q-dist.eps}  \\caption{Mass ratio distribution for Am binary systems. The solid-line histogram is the distribution based on eclipsing binaries in the current work, while the solid grey histogram is that presented by \\cite{2010A&A...524A..14B} based on the spectroscopic sample. The dashed histogram is the estimated eclipsing binary distribution after making allowance for the WASP detection probabilities.}  \\label{q-dist} \\end{figure}  As discussed in Sect.~\\ref{ecl_prob}, eccentric systems may not always show two eclipses when both stars are similar. For example Renson\\,42906 (HD\\,151604; V916 Her) is an eccentric ($e = 0.566$) system with mass ratio close to unity \\citep{2007MNRAS.380.1064C}, but the WASP data only shows one eclipse per orbit and value of $q=0.12$. Thus, these systems will appear to have anomalously low $q$ values, further adding to the uncertainty in the mass-ratio distribution.  \\cite{2010A&A...524A..14B} concluded that the mass-ratio distribution showed hints of a double-peaked distribution, with peaks at $q \\sim 0.3$ and $q \\sim 1$. The mass-ratio distribution based on the estimated properties from the WASP light curves for detached systems is noticeably different (Fig.\\,\\ref{q-dist}). The estimated WASP mass-ratio distribution shows a broad peak near unity, with a deficit around $q \\simeq 0.3$. However, the WASP detection probability varies with companion size (Sect.\\,\\ref{det_prob}). Assuming that companions with masses around 0.5\\,{$M_\\sun$} correspond to the transition between the two system scenarios discussed earlier, there would be an increase in the number of low $q$ systems relative to the high $q$ systems. The distribution would become flatter,  similar to that found by \\cite{2010A&A...524A..14B}, who noted that a flat mass-ratio distribution also appeared to be a good fit. While there are genuinely low $q$ systems (e.g. Renson 3750), the apparent excess of such systems may not be real, since some of these may be pairs of similar stars with a true period equal to twice the assumed one (as noted with a dagger in Table~\\ref{table:Am Binaries}). Hence, since our mass-ratio estimates are based on photometry alone, radial velocity studies are required to determine spectroscopic mass-ratios of the whole sample, before any firmer conclusions can be drawn."
    },
    "1402/1402.4524.txt": {
        "abstract": "We present {\\it Herschel}/PACS 100 and 160 $\\mu$m integrated photometry for the 323 galaxies in the {\\it Herschel} Reference Survey (HRS), a K-band-, volume-limited sample of galaxies in the local Universe. Once combined with the {\\it Herschel}/SPIRE observations already available, these data make the HRS the largest representative sample of nearby galaxies with homogeneous coverage across the 100-500 $\\mu$m wavelength range. In this paper, we take advantage of this unique dataset to investigate the properties and shape of the far-infrared/sub-millimeter spectral energy distribution in nearby galaxies. We show that, in the stellar mass range covered by the HRS (8\\ls$\\log(M_{*}/M_{\\odot})$\\ls12), the far-infrared/sub-millimeter colours are inconsistent with a single modified black-body having the same dust emissivity index $\\beta$ for all galaxies. In particular, either $\\beta$ decreases, or multiple temperature components are needed, when moving from metal-rich/gas-poor to metal-poor/gas-rich galaxies. We thus investigate how the dust temperature and mass obtained from a single modified black-body depend on the assumptions made on $\\beta$. We show that, while the correlations between dust temperature, galaxy structure and star formation rate are strongly model dependent, the dust mass scaling relations are much more reliable, and variations of $\\beta$ only change the strength of the observed trends. %Our results suggests that far-infrared/sub-millimeter colours may not always be a good proxy for the average %temperature of the cold dust reservoir in galaxies. ",
        "introduction": "It is now well established that approximately half of the radiative energy produced by galaxies is absorbed by dust grains and re-emitted in the infrared regime \\citep{hauser01,bosellised,dole06,dale07,burga2013}. Thus, observations in the $\\sim$10-1000 $\\mu$m wavelength range provide us with a unique opportunity not only to quantify half of the bolometric luminosity of galaxies, but also to characterise the properties of cosmic dust. Moreover, since dust grains are crucial for the star formation cycle \\citep{hollenbach71}, such information can give us important insights into the physical processes regulating galaxy evolution (e.g., \\citealp{dunne11}).  Unfortunately, despite its paramount importance, we are still missing a complete and coherent picture of dust properties in galaxies across the Hubble sequence, and of the exact role played by grains in regulating star formation \\citep{mckee10}. Indeed, we know very little about the dust composition in galaxies outside our own Local Group \\citep{draine07b,compiegne11} and if/how it is regulated by the physical conditions experienced by grains in the inter-stellar medium (ISM). Hence, our estimates of dust masses in galaxies are still highly uncertain \\citep{fink99,dupac03,gordon10,paradis10,plankbeta}.  Luckily, the last decade has seen the start of a golden age for observational far-infrared (FIR) and sub-millimeter (submm) astronomy, providing a new boost to the refinement of theoretical dust models \\citep{meny07,draine07b,hoang2010,compiegne11,steinacker13}. In particular, the {\\it Spitzer} \\citep{spitzer}, and more recently {\\it Herschel} \\citep{pilbratt10} and {\\it Planck} \\citep{planck} space telescopes are finally gathering a wealth of information on the dust emission from thousands of galaxies up to $z\\sim$2. Particularly important for a proper characterisation of dust in galaxies is the radiation emitted at wavelengths \\gs100-200 $\\mu$m. In this regime, the integrated emission from galaxies originates predominantly from dust in thermal equilibrium, heated by the diffuse interstellar radiation field (ISRF), which represents the bulk of the dust mass in a galaxy (e.g., \\citealp{sodrosky89,sauvage92,calzetti95,walterbos,bendo10,boquien11,bendo12b}). Thus, by characterising the dust emission in the \\gs100 $\\mu$m wavelength domain, we have a unique opportunity to provide strong constraints to theoretical models, and to refine our census of the dust budget in galaxies.  The first natural step in this direction is to quantify how the shape of the dust spectral energy distribution (SED) varies with galaxy properties across a wide range of morphological type, star formation activity, cold gas mass and metal content. This is necessary to determine if the amount of radiation emitted at each wavelength is simply regulated by the intensity of the ISRF responsible for the dust heating, or whether it retains an imprint of the chemical composition of the grains. Indeed, only after a careful characterisation of the physical parameters regulating the dust SED, will it be possible to properly convert observables into physical quantities such as dust temperatures and dust masses.  Many recent works \\citep{gordon10,skibba11,davies11,plankbeta,galametz12,auld13} have shown that, above $\\sim$100 $\\mu$m, the dust SED is very well approximated by a simple modified black-body (but see also \\citealp{bendo12b}): \\begin{equation} F_{\\nu} = \\frac{M_{dust}}{D^{2}} \\kappa_{\\nu_{0}}\\left(\\frac{\\nu}{\\nu_{0}}\\right)^{\\beta} B_{\\nu}(T) \\label{eq1} \\end{equation} where $F_{\\nu}$ is the flux density emitted at the frequency $\\nu$, $\\kappa_{\\nu_{0}}$ is the dust mass absorption coefficient at the frequency $\\nu_{0}$, $\\beta$ gives its variation as a function of frequency, $D$ is the galaxy distance and $B_{\\nu}(T)$ is the Planck function. Mounting evidence is emerging that $\\beta$ is not the same in all galaxies (e.g., \\citealp{remy13}), and may also vary within galaxies (e.g., \\citealp{galametz12,smith12b}).  Modified black-bodies are simple models and cannot properly reproduce real dust properties (e.g., \\citealp{draine07b,shetty09b,bernard10}). Several dust components at various temperatures contribute to the total emission along the lines-of-sight. This implies the presence of temperature mixing that can cause variations of the infrared slope, and thus in the apparent emissivity index $\\beta$. Nevertheless, parameterization of the dust SEDs through modified black-body fitting is a powerful tool to help understand variations of dust properties with other galaxy characteristics, especially in case of sparse sampling of the FIR/sub-mm wavelength range (e.g., high-redshift galaxies \\citealp{magdis2011,symenodis2013}). Therefore, it is extremely important to determine in which cases a single modified back-body can be used, and how temperature and dust mass estimates are affected by the assumptions made on $\\beta$.  In order to ascertain the dust properties of galaxies in the local Universe, and to provide new constraints to theoretical models, we have carried out the {\\it Herschel} Reference Survey (HRS, \\citealp{HRS}), a {\\it Herschel} guaranteed time project focused on the study of the interplay between dust, gas and star formation in a statistically significant sample of $\\sim$300 galaxies spanning a wide range of morphologies, stellar masses (8\\ls log(M$_{*}$/M$_{\\odot})$\\ls 12), cold gas contents (-3\\ls log(M$_{HI}$/M$_{*}$)\\ls 1), metallicities (8.2\\ls 12+log(O/H) \\ls 8.9), and specific star formation rates (-12 \\ls log(SFR/M$_{*}$)\\ls -9). The combination of {\\it Herschel}/SPIRE \\citep{spire} observations with the multi-wavelength dataset we have been assembling \\citep{ciesla12,hrsgalex,boselli13,hughes13}, has already allowed us to have a first glimpse at how the dust content and shape of the dust SED vary with internal galaxy properties \\citep{boselli10,boselli12,cortese12}. In particular, \\cite{boselli10,boselli12} have shown that the slope of the dust SED in the 200-500 $\\mu$m interval decreases from $\\beta\\sim$2 to $\\beta\\sim$1 when moving from metal-rich to metal-poor galaxies. However, our analyses have so far been limited by the lack of data in the $\\sim$100-200 $\\mu$m wavelength range for the entire sample.  Thus, in this paper we present integrated {\\it Herschel}/PACS \\citep{pacs} 100 and 160 $\\mu$m flux densities for all the HRS sample and take advantage of our multiwavelength dataset to perform a first analysis of the properties of the dust SED across our entire sample. Corresponding to the peak of the dust SED, the 100-200 $\\mu$m wavelength interval is crucial not only to properly quantify the shape of the SED, but also to accurately determine the average dust temperature and total dust mass in galaxies. These data make the HRS the largest representative sample of nearby galaxies with homogeneous coverage across the $\\sim$100-500 $\\mu$m wavelength range. In addition to releasing our dataset to the community, our primary goals are 1) to investigate how the shape of the dust SED varies with internal galaxy properties, and 2) to determine whether the integrated dust SED of HRS galaxies can always be reduced to a single modified black-body with a constant value of $\\beta$ and, if not, what are the possible biases introduced by this assumption. The results of SED fitting with the dust models of \\cite{draine07} will be presented in a forthcoming paper (Ciesla et al., submitted.).  This paper is organized as follows. In Sect. 2 we describe the {\\it Herschel} observations, data reduction, flux density estimates and comparison with the literature. In Sect. 3 we use the PACS and SPIRE colours to investigate how the shape of the dust SED varies with internal galaxy properties. In Sec. 4, we show how the dust temperature and mass obtained from fitting a single modified black-body to the {\\it Herschel} data depend on the assumptions made on $\\beta$. Finally, the summary and implications of our results are presented in Sec. 5.    \\begin{figure*} \\centering \\includegraphics[width=17.5cm, angle=0]{images.epsi} \\caption{Comparison of the quality of our PACS images with the Sloan Digital Sky Survey optical and SPIRE 250 $\\mu$m images. We show three types of objects: an early-type with dust lanes (top row), an unperturbed late-type spiral and an un-detected elliptical and its spiral companion. The size of the PACS and SPIRE beams is shown in the bottom left corner of each panel.} \\label{images} \\end{figure*} ",
        "conclusions": "In this paper we presented PACS 100 and 160 $\\mu$m integrated photometry for the {\\it Herschel} Reference Survey. We have combined these data with SPIRE observations to investigate how the shape of dust SED varies across the Hubble sequence. Being the largest representative sample of nearby galaxies with homogeneous coverage in the 100-500 $\\mu$m wavelength domain, the HRS is ideal to quantify if and how dust emission varies across the local galaxy population. Our main results are as follows.  \\begin{itemize}  \\item The shape of the dust SED is not well described by a single modified black-body having just the dust temperature as a free parameter. Instead, there is a clear need to vary the dependence of the dust emissivity ($\\beta$) on wavelength, or to invoke multiple temperature components in order to reproduce the colours observed in our sample. This is particularly important as the HRS does not include very metal-poor dwarf galaxies, for which we already knew that the dust SED is significantly different from the one of metal-rich, massive galaxies \\citep{galliano2005,galliano11,engelbracht2008,galametz2009,remy13}. Our results suggest that the difference in FIR/sub-mm colours between giant and dwarf galaxies \\citep{draine07b} may not be the result of a dramatic transition in dust properties, but just the consequence of the gradual variation that we observe as a function of metal and gas content.  \\item The variation in the slope of the dust SED strongly affects dust temperature estimates from single modified black-bodies fits. In particular, the correlations between galaxy properties and dust temperatures strongly depend on the assumptions made on $\\beta$: i.e., trends can disappear or even reverse. Conversely, dust mass estimates are more robust, and variations in $\\beta$ do not produce the same dramatic inversion of some correlations observed for the dust temperature.  \\item We confirm that the temperature of a single modified black-body is mainly related to specific star formation rate, while $\\beta$ varies more with the degree of metal enrichment of the ISM.   \\end{itemize}  The results presented in this paper may appear in contradiction with several recent works showing that the dust SED is very well reproduced by a simple modified black-body with $\\beta=$2 \\citep{davies11,auld13}. However, all these works were focused on massive, metal-rich and relative gas-poor galaxies, for which we also find that a constant value of $\\beta$ provides a good fit to our data. It is when we move to the gas-rich/metal-poor regime that the shape of the SED starts to change \\citep{boselli10,boselli12,remy13}.  Our findings overall reinforce the results already presented in \\cite{boselli10,boselli12}. However, it is important to note that the discovery of a clear variation in the shape of the SED across the HRS has only been possible thanks to the large wavelength coverage obtained by combining both PACS and SPIRE data. Indeed, with SPIRE or PACS data only, it would be not only much more difficult to show under which conditions a simple modified black-body approach does not work, but it would also be nearly impossible to quantify how model assumptions can affect the correlation of dust temperature with star formation, galaxy structure and chemical enrichment.   %Despite this, {\\it Herschel} colours still represent an ideal tool to quantify %how the dust SED varies with other characteristics of the ISM in galaxies. %{\\bf Indeed, once we take into account that PACS and SPIRE colours are related to different %dust properties, they can be used as proxies for $T$ and $\\beta$ without being %affected by any degeneracy hidden in the SED fitting technique. %Thus, whenever possible, we suggest using colours to carefully test the reliability %of relations obtained as results of a $\\chi^{2}$ minimizations of a single %modified black-body.   %% Finally, the need of varying $\\beta$ even for a sample like the HRS has also profound implications %% on our ability to quantify dust masses in galaxies, as a variation in the value %% of the dust emissivity used to normalise the modified black-body is required \\citep{bianchi13}. %% Of course, this automatically affects the use of far-infrared/sub-millimeter %% observations to constrain the amount of cold gas \\citep{james02,eales2010} and/or the value of the %% CO-to-H$_{2}$ conversion factor in nearby \\citep{guelin93,boselligdust,leroy2011} and high-redshift galaxies \\citep{magdis2011}. %% For example, the simple exercise described in the previous section shows that %% dust masses could be easily underestimated by a factor of $\\sim$2, when moving from %% high to low metallicity environments. This would automatically lead to an underestimate %% of the variation of the CO-to-H$_{2}$ conversion factor with metallicity. %% Unfortunately, no dust models are currently able to predict the evolution %% of the dust emissivity with $\\beta$ across the range of metallicities and gas content %% covered by the HRS. This is an issue that, thanks to the complete coverage now available %% for our sample, we plan to investigate in future work. %% In the meantime, we recommend caution when adopting the same dust model for %% quantifying the dust properties and dust mass for samples of galaxies covering %% a wide range in metallicity, cold gas content and/or specific star formation rates."
    },
    "1402/1402.6358_arXiv.txt": {
        "abstract": "We present a quantitative analysis of the low-resolution ($\\sim$4.5~\\AA) spectra of 12 late-B and early-A blue supergiants (BSGs) in the metal-poor dwarf galaxy NGC~3109. A modified method of analysis is presented which does not require use of the Balmer jump as an independent T$_{eff}$ indicator, as used in previous studies. We determine stellar effective temperatures, gravities, metallicities, reddening, and luminosities, and combine our sample with the early-B type BSGs analyzed by \\citet{E07} to derive the distance to NGC 3109 using the Flux-weighted Gravity-Luminosity Relation (FGLR). Using primarily Fe-group elements, we find an average metallicity of [\\={Z}] = -0.67 $\\pm$ 0.13, and no evidence of a metallicity gradient in the galaxy. Our metallicities are higher than those found by \\citet{E07} based on the oxygen abundances of early-B supergiants ([\\={Z}] = $-$0.93 $\\pm$ 0.07), suggesting a low $\\alpha$/Fe ratio for the galaxy. We adjust the position of NGC 3109 on the BSG-determined galaxy mass-metallicity relation accordingly and compare it to metallicity studies of HII regions in star-forming galaxies. We derive an FGLR distance modulus of 25.55 $\\pm$ 0.09 (1.27 Mpc) that compares well with Cepheid and tip of the red giant branch (TRGB) distances. The FGLR itself is consistent with those found in other galaxies, demonstrating the reliability of this method as a measure of extragalactic distances. ",
        "introduction": "The extreme brightness of blue supergiants (BSGs), a short post-main sequence evolutionary stage of 12~M$_{\\odot}$ to 40~M$_{\\odot}$ stars, makes it possible to obtain resolved spectra of individual BSGs out to 10 Mpc with current instrumentation. As such, BSGs are ideal tools to obtain crucial information about the chemical composition of nearby galaxies and provide insight to their chemical evolution \\citep{K08,K12}. Often galaxy metallicities are studied through the spectroscopy of HII regions, which has been widely applied to examine radial abundance gradients of spiral galaxies \\citep{VC92, Z94, P04} and the galaxy mass-metallicity relation \\citep{L79, T04, A13}. However, this approach is limited by its reliance on empirical ``strong line'' analysis methods, which have been shown to yield significantly different absolute metallicities depending on what calibration is used \\citep{Ke08, B09}. Even in cases where metallicities can be measured more directly using the weak auroral lines, HII region studies might be affected by systematic uncertainties difficult to assess, such as oxygen depletion on to dust grains and a possible detection bias towards lower abundances at high metallicities \\citep{ZB12}. BSGs thus provide a valuable independent measure of galaxy metallicity.  In addition, it has been shown that BSGs can be used as distance indicators through the Flux-weighted Gravity-Luminosity Relation (FGLR, Kudritzki et al. 2003, 2008). This relation correlates stellar gravity and effective temperature, which can be derived from the stellar spectrum, to the absolute bolometric magnitude. The FGLR is advantageous in that it is free of uncertainties caused by interstellar reddening, since the reddening is determined during the spectral analysis. In addition, a potential metallicity dependence of the FGLR, if present, can be accounted for since metallicity is also determined independently by the analysis of each object. This is especially valuable in light of recent efforts to establish the Hubble constant \\emph{H$_{0}$} to an accuracy better than 5\\%, which would greatly constrain cosmological parameters without having to invoke assumptions about the geometry of the universe \\citep{KU12, Riess11}. FGLR distances found for WLM \\citep{U08}, M33 \\citep{U09}, and M81 \\citep{K12} have been found to be consistent with distances determined by other methods, demonstrating the reliability of the method.  We conduct a spectroscopic study of 12 late-B and early-A type BSGs in NGC~3109, a Magellanic SBm galaxy at the edge of the Local Group \\citep{dV91,Bergh99}. With M$_{V}$ = -14.9 \\citep{Mcc12}, the galaxy is the most luminous member of the NGC~3109 group, which according to \\citet{Tully06} is the ``nearest distinct structure of multiple galaxies to the Local Group''. Recent work by \\citet{ST13} and \\citet{B13} indicate that the members of this group form a $\\sim$1070 kpc filamentary structure created by tidal interaction or filamentary accretion. Our purpose is two-fold: to determine the FGLR and calculate a corresponding distance to NGC~3109, and to evaluate its average metallicity using multiple Fe-group element species. An analysis of eight early B-type BSGs by Evans et al. (2007, hereafter E07) found the oxygen abundance of NGC~3109 to be approximately 1/10 of solar, a result consistent with HII regions studies using auroral lines \\citep{Le03, Le03b, P07}. If these values reflect the overall metallicity, NGC~3109 will be the lowest metallicity object for which an FGLR has been constructed, allowing us to investigate the metallicity dependence of the relation and compare with stellar evolution theory. Since we fit metal lines of multiple elements, in particular Fe-group elements, our metallicities will more closely resemble the overall stellar metallicities and are not restricted to oxygen as a proxy of stellar metallicity as in the case of the HII region studies and the work of E07. On the other hand, a comparison with the oxygen abundances obtained in these studies will provide insight to the $\\alpha$/Fe abundance ratio, a key parameter in constraining the chemical evolution and star formation history of the galaxy. ",
        "conclusions": "We analyze low-resolution ($\\sim$4.5~\\AA) spectra of 12 late-B and early-A BSGs in NGC~3109, obtaining the effective temperatures, gravities, metallicities, reddening, and luminosities of these objects. We use a modified method of analysis that does not use the Balmer jump to break the temperature-metallicity degeneracy. Instead, we employ a $\\chi^2$-based approach that takes advantage of the fact that spectral lines from different atomic species and from different excited levels react differently to changes of temperature and metallicity and use this to constrain stellar parameters. A test analysis of SMC spectra is found to produce parameters consistent with high-resolution analyses, attesting to the accuracy of our technique. A disadvantage of this method is that we must make assumptions regarding the He abundances of our objects, which we cannot determine from our analysis alone. Thus we consider two sets of model atmospheres, one assuming a ``normal'' He abundance (based on averages from BSG studies in the MW) and another assuming an enhanced He abundance, and take the average metallicity and FGLR distance as our final results for NGC 3109. Fortunately, the adopted He abundance only has a small impact on the derived metallicities, with the temperatures, gravities, reddening, and luminosities being very similar for both sets of models.  From our sample, we find the average Fe-group metallicity of NGC~3109 to be [\\={Z}] = $-$0.63 $\\pm$ 0.13 in the normal He abundance case and [\\={Z}] = $-$0.71 $\\pm$ 0.14 in the enhanced He abundance case, resulting in our adopted metallicity of  [\\={Z}] = $-$0.67 $\\pm$ 0.13. Even in the enhanced He abundance case, this result is higher than the oxygen-based metallicity obtained by \\citet{E07}, who find an average value of [\\={Z}] = $-$0.93 $\\pm$ 0.07 based on an analysis of 8 early-B BSGs. This may indicate a sub-solar $\\alpha$/Fe ratio in the galaxy. We adjust the position of NGC~3109 on the BSG-based galaxy mass-metallicity relation presented in \\citet{K12},  and find that the relation compares well (in a qualitative sense) to the recent mass-metallicity relation of \\citet{A13} based on auroral line metallicity measurements of star-forming galaxies in SDSS. Interestingly, the metallicities of the \\citet{A13} relation appear to be systematically higher than those in the BSG relation, an inconsistency which requires further investigation.  We combine our results with the BSGs analyzed by \\citet{E07} to determine the Flux-weighted Gravity-Luminosity Relation (FGLR) of NGC~3109. We find the FGLR to be almost identical to those found in other galaxies, demonstrating the consistency of the relation across a wide range of galaxy masses and metallicities. We obtain an FGLR distance modulus of $\\mu$ = 25.55 $\\pm$ 0.09 (1.27 Mpc) that is effectively independent of the adopted stellar He abundances. This result is in good agreement with distances found in Cepheid variable and TRGB studies, serving as an independent confirmation of these values. The consistency between our FGLR distance and the Cepheid distances of \\citet{Pi06} and \\citet{So06} suggests that the Cepheid period-luminosity (PL) relation is not strongly affected by metallicity, since these studies adopt PL slopes derived for higher metallicities ([Z]~=~$-$0.3). This study is an additional example of the great value of blue supergiants as independent distance and metallicity indicators in nearby galaxies."
    },
    "1402/1402.6672_arXiv.txt": {
        "abstract": "We perform population synthesis simulations for Population III (Pop III) coalescing compact binary which merge within the age of the universe. We found that the typical mass of Pop III binary black holes (BH-BHs) is $\\sim30~\\msun$ so that the inspiral chirp signal of gravitational waves can be detected up to z=0.28 by KAGRA, Adv. LIGO, Adv. Virgo and GEO network. Our simulations suggest that the detection rate of the coalescing Pop III BH-BHs is $140(68)\\ {\\rm events/yr ~(SFR_p/(10^{-2.5}\\msun/yr/Mpc^3))\\cdot Err_{sys}}$ for the flat (Salpeter) initial mass function (IMF), respectively, where $\\rm SFR_{p}$ and $\\rm Err_{sys}$ are the peak value of the Pop III star formation rate and the possible systematic errors due to the assumptions in Pop III population synthesis, respectively. $\\rm Err_{sys}=1$ correspond to conventional parameters for Pop I stars. From the observation of the chirp signal of the coalescing Pop III BH-BHs, we can determine both the mass and the redshift of the binary for the cosmological parameters determined by Planck satellite. Our simulations suggest that the cumulative redshift distribution of the coalescing Pop III BH-BHs depends almost only on the cosmological parameters. We might be able to confirm the existence of Pop III massive stars of mass $\\sim 30~\\msun$ by the detections of gravitational waves if the merger rate of the Pop III massive BH-BHs dominates that of Pop I BH-BHs. ",
        "introduction": "Gravitational-wave astronomy with KAGRA\\footnote{http://gwcenter.icrr.u-tokyo.ac.jp/en/}, Adv. LIGO\\footnote{http://www.ligo.caltech.edu/},  Adv. Virgo\\footnote{http://www.ego-gw.it/index.aspx/}, and GEO\\footnote{http://www.geo600.org/}  will reveal the formation and evolution of binaries through the observed merger rates of compact binaries, such as binary neutron stars (NS-NSs), neutron star -- black hole binaries (NS-BHs), and binary black holes (BH-BHs).  For this gravitational wave astronomy, estimates of the merger rate of compact binaries play key roles to develop observational strategy and to translate the observed merger rates into the binary formation and evolution processes.  There are two methods to estimate the merger rate of compact binaries.   One is to use observational facts such as the observed NS-NSs whose coalescence time due to the emission of gravitational waves is less than the age of the universe. Taking into account the observation time,  the sensitivity of the radio telescope, the luminosity function of pulsars and the beaming factor so on, the probability distribution function of the merger rate can be found. For example, \\citet{Kalogera2004b} found that the event rate of the coalescing NS-NSs is in the range from $10^{-5}\\ {\\rm events}\\ {\\rm yr}^{-1}\\ {\\rm galaxy}^{-1}$ to $4\\times 10^{-4}\\ {\\rm events}\\ {\\rm yr}^{-1}\\ {\\rm galaxy}^{-1}$ at the $99\\ \\%$ confidence level (see their Fig.~2)\\footnote{Note here that there are errors in \\citet{Kalogera2004a} so that the rates in \\citet{Kalogera2004b}  are the correct ones.}.   The merger rate of NS-NSs can be restricted by  the rate of the observed Type Ib and Ic supernovae, supposing that the formation of NS-NSs really starts from the massive binary zero age main sequence~(ZAMS) stars. This is because the formation of the second neutron star should occur in association with Type Ib and Ic supernovae in which the H-rich envelope and the He-layer are lost, respectively, otherwise the binary disrupts due to the sudden large mass loss at the supernova explosion\\footnote{If more than half of the total mass is suddenly lost at the supernova explosion, the binary disrupts.}. Under the assumption of the equality of the formation rate to the merger rate, the merger rate of the NS-NSs is limited by the Type Ib and Ic supernova rate of $\\sim 10^{-3}\\ {\\rm events}\\ {\\rm yr}^{-1}\\ {\\rm galaxy}^{-1}$ \\citep{Cappellaro1997, Cappellaro1999,Li2011}.  Therefore  the maximum rate of $4\\times 10^{-4}\\ {\\rm events}\\ {\\rm yr}^{-1}\\ {\\rm galaxy}^{-1}$ by \\citet{Kalogera2004b} implies that $\\sim 40\\ \\%$ of the Type Ib and Ic supernovae is associated with the formation of NS-NSs with the coalescence time less than the age of the universe. This percentage seems to be too large.  We also note here that under the assumption that central engine of short gamma ray bursts are coalescing binary neutron stars, one  can use the observations of short gamma ray bursts to estimate the coalescing rate (\\citet{coward12} ~and references cited there).    The dynamical interaction in a globular cluster is another route to the formation of NS-NSs since there exists PSR2127+11C, which is contained in a NS-NS system, in the globular cluster M15 \\citep{Prince1991}. The age of the globular cluster is $\\sim 10^{10}\\ {\\rm yr}$ so that all the massive stars ended their life and the young pulsars do not exist.  The coalescence time of PSR2127+11C is $\\sim 2 \\times 10^{8}\\ {\\rm yr}$ which is much smaller than the age of the globular cluster so that it was formed most likely by three body interactions such as the collision of neutron star -- white dwarf binary or neutron star -- dwarf star binary with a single neutron star. Since there exists only one NS-NS observed in the globular cluster, it is difficult to estimate the merger rate. Theoretical simulation is the only method at present~\\citep{Grindlay2006, Ivanova2008}.    The another method to estimate the merger rate of compact binaries is to use theoretical computation based on the hypothetical assumptions of binary formation and evolution. For NS-BHs and BH-BHs, in particular, there exists no observations so that we can only use theoretical estimates. The merger rates of compact binaries of Population I~(Pop I) stars were estimated  by many authors~\\citep{Belczynski2002, Belczynski2007, Belczynski2012}.  \\citet{Dominik2012} computed the merger rates for the progenitor stars  of metallicity $Z=0.1\\ Z_\\odot$ and found that the number of the coalescing BH-BHs increases compared with that for $Z=Z_\\odot$. \\citet{Dominik2013} adopted a certain model of the star formation rate and the chemical evolution of the metallicity $Z$ to compute the cumulative redshift distribution of the coalescing compact binaries.    In this paper, we focus on the compact binary merger originated from Population III stars (Pop III stars) as gravitational wave sources. Pop III stars are the first stars after the Big Bang which are formed from metal-free gas~\\citep{Omukai1998,  Bromm2002,  Abel2002, Yoshida2008, Greif2012}. The simulations of a rotating primordial gas cloud suggest that the formation of Pop III star binaries and multiple star systems are frequent \\citep{Machida2008a, Stacy2010}. The main differences of our work from \\cite{Dominik2012,Dominik2013} are the following two: (1) we focus on metal-free Pop III stars and (2) consider the star formation history including the transition to metal-enriched stars (see \\S4). The observed merger rate will be the sum of our work and \\citet{Dominik2013}.      There are at least three differences between Pop III and Pop I compact binaries. First of all, Pop III stars are more massive than Pop I stars \\citep{McKee2008, Hosokawa2011, Hosokawa2012, Stacy2012} with mass $10-100\\ {\\rm M}_\\odot$ so that  Pop III star binaries probably evolve into BH-BHs.  Secondly, since the typical formation time of Pop III stars is at $z\\sim 10$, even if the coalescence time is comparable to the age of the universe, they merge at present and contribute to the sources of gravitational waves for KAGRA, Adv. LIGO, Adv. Virgo, and GEO network. Therefore, if Pop III NS-NSs were formed, they might merge at present so that their rate is free from the constraint of the observed NS-NSs as well as Type Ib and Ic supernova rate discussed in the previous paragraphs. Thirdly, Pop III black holes are expected to be more massive than Pop I black holes due to less mass loss so that the resulting gravitational waves are easier to detect, since the detectable distance is proportional to 5/6 power of the chirp mass ($M_{\\rm chirp}$) of a binary defined by $ M_{\\rm chirp} = (M_1M_2)^{3/5}/(M_1 + M_2)^{1/5}$ (Peters 1964 and Peters \\& Mathews 1963, see also Sathyaprakash \\& Schutz(2009)), where $M_1$ and $M_2$ are the mass of each compact object.    The idea of Pop III compact binaries as gravitational-wave sources has been considered by \\cite{Belczynski2004}, \\cite{Kulczycki2006} and \\cite{Kowalska2012}. However they considered very massive Pop III stars with mass over hundred solar masses. Recent study shows that the typical mass of Pop III stars is set to 10--100$\\rm{M_{\\odot}}$ by the stellar radiation feedback on the accretion flow~\\citep{Hosokawa2011, Hosokawa2012}. Therefore, in this paper, we calculate $10^6$ Pop III binary evolutions with the mass range of 10--100$\\rm{M_{\\odot}}$ to estimate merger rates and mass distribution of Pop III compact binaries. In order to calculate Pop III binary evolutions, we upgrade Hurley's binary population synthesis code \\citep{Hurley2002} for the Pop I star to Pop III star case.  This paper is organized as follows. We describe Pop III single star evolution in \\S \\ref{sec:single}, the method to calculate Pop III binary star evolutions in \\S \\ref{sec:binary}, numerical calculation methods in \\S \\ref{sec:initial condition}. In \\S 3, we present the results of simulations and argue properties of Pop III compact binaries. We compare Pop III compact binary mergers with Pop I compact binary mergers in \\S \\ref{sec:comparison models}. In \\S 4, we describe the Pop III compact binary merger rates. \\S 5 is devoted to the discussions. In Appendix, we show the details of our numerical methods, the comparison of our results with Hurley's ones and the convergence check of our simulations. We adopt the cosmological parameters of $(\\Omega_{\\Lambda}, \\Omega_{\\rm m}) = (0.6825, 0.3175)$ and the Hubble parameter of $H_0 = 67.11\\ {\\rm km}\\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1}$~\\citep{Planck2013}. Those who are not interested in the details of the methods in  numerical simulations  can skip \\S 2. ",
        "conclusions": "In this paper, we do not include the magnetic braking since Pop III stars are formed from the primordial no metal gas. However, there are discussions against this assumption. Among them, the enhancement of the magnetic field during the star formation has been well studied using the numerical simulations \\citep[e.g.,][]{Maki2004,Machida2008b,Machida2013}. According to their results, the turbulent motions driven in the galaxy formation might increase the magnetic field up to $< 10^{-6}$ G \\citep{Schleicher2010,Sur2010,Schober2012,Turk2012}, which is similar value to that in molecular clouds of the galaxy. If this is the case, we should include the effect of magnetic braking, which will give rise $\\rm Err_{sys}\\neq 1$.  \\cite{Dominik2012} discussed the metallicity dependence of the compact binary merger by focusing on Pop I stars with metallicities ${\\rm Z_{\\odot}}$ and $0.1\\ {\\rm Z_{\\odot}}$. Here, we compare our results for the coalescing Pop III BH-BHs with those in \\cite{Dominik2012} and discuss the implications of our results~(see Table~\\ref{dominik}). In Table~\\ref{dominik}, the second and third columns show the results of \\cite{Dominik2012} for metallicity $\\rm Z_\\odot$ and $\\rm 0.1\\ Z_\\odot$ stars. Here, Models A and B correspond to the standard case of submodels A and B in \\cite{Dominik2012}. The last column show our results for Pop III binaries, where we take the fiducial parameter values: $\\rm Err_{sys}= 1$ and $\\rm SFR_p=10^{-2.5}\\ \\msun\\ yr^{-1}\\ Mpc^{-3}$. As in Table \\ref{dominik}, \\cite{Dominik2012} suggested that for Pop I stars with $Z = {\\rm Z_{\\odot}}$, the merger rate of the BH-BHs becomes $8.2\\ (1.9)\\times 10^{-8}\\ {\\rm events\\ yr^{-1}\\ Mpc^{-3}}$ in their Model A~(Model B), respectively, while for Pop I stars with $Z = 0.1\\ {\\rm Z_{\\odot}}$, it becomes $7.33\\ (1.36) \\times 10^{-7}\\ {\\rm events\\ yr^{-1}\\ Mpc^{-3}}$ in their Model A~(Model B), respectively. Here we assume that the galaxy with metallicity $Z=0.1\\ {\\rm Z}_\\odot$ has the hypothetical number density $\\sim 10^{-2}\\ \\rm Mpc^{-3}$ in order to compare the results. On the other hand, for our Pop III BH-BHs, the expected merger rate is $2.5\\ (1.2)\\times 10^{-8}\\ {\\rm events\\ yr^{-1}\\ Mpc^{-3}}$ for the flat (Salpeter) IMF, respectively. This means that the merger rate of the Pop III BH-BHs is smaller than that of Pop I. However, for Pop I BH-BHs with metallicity $Z={\\rm Z}_\\odot\\ (0.1\\ {\\rm Z}_\\odot)$, the typical chirp mass is 6.7~(13.2) $\\msun$, while for Pop III ones it is $\\sim 30\\ \\msun$ as we showed in Fig.~7. Since the detection range of merger events increases in proportion to $M_{\\rm chirp}^{5/6}$, the detection rate increases in proportion to $M_{\\rm chirp}^{5/2}$. Therefore, the detectable event rate of the Pop III BH-BHs is 13~(6) times larger for the flat (Salpeter) IMF than that of Pop I with $Z={\\rm Z}_\\odot$ in Model A, while it is 0.26~(0.13) times larger than that of Pop I with $Z=0.1\\ {\\rm Z}_\\odot$, if the galaxy consists of stars with $Z=0.1\\ {\\rm Z}_\\odot$. In Model B, these numbers become 56~(26) for $Z={\\rm Z}_\\odot$ case and 1.4~(0.7) for $Z=0.1\\ {\\rm Z}_\\odot$ case, respectively. Thus, for the fiducial parameters of Pop III binaries, the contribution of Pop III BH-BHs is comparable to or larger than that of Pop I with $Z={\\rm Z}_\\odot$ and $0.1\\ {\\rm Z}_\\odot$, since the major part of a galaxy does not necessarily consist of Pop I stars with $Z=0.1\\ {\\rm Z}_\\odot$.  \\begin{table*} \\caption[]{The comparison of the merger rate density of the BH-BHs and typical chirp mass between previous studies and our study. The second and third columns show the results of \\cite{Dominik2012} for metallicity $\\rm Z_\\odot$ and $\\rm 0.1\\ Z_\\odot$ stars. Here, Models A and B correspond to the standard case of submodels A and B in \\cite{Dominik2012}. The last column show our results for Pop III binaries. Here, we take the fiducial parameter values: $\\rm Err_{sys}= 1$ and $\\rm SFR_p=10^{-2.5}\\ \\msun\\ yr^{-1}\\ Mpc^{-3}$.} \\label{dominik} \\begin{tabular}{l c c c} % \\hline & Z$_\\odot$ & 0.1 Z$_\\odot$  & Pop III\\\\%\\midrule \\hline Model A [$10^{-8}\\ {\\rm events\\ yr^{-1}\\ Mpc^{-3}}$]             & 8.2             & 73.3                 & 2.5 (flat)\\\\ Model B [$10^{-8}\\ {\\rm events\\ yr^{-1}\\ Mpc^{-3}}$]             & 1.9             & 13.6                 & 1.2 (Salpeter)\\\\ chirp mass [$\\msun$]                                                           & 6.7              &13.2                  & 30\\\\ \\hline \\end{tabular} \\end{table*}  While \\cite{Dominik2012} did not take account of the evolution of the star formation rate, \\cite{Dominik2013} adopted a certain model of the star formation rate~\\citep[Eq.~1 of][]{Dominik2013} and the metallicity $Z$ evolution~\\citep[Eqs. 3 to 5 of][]{Dominik2013} to compute the cumulative redshift distribution of the coalescing compact binaries. They also took into account the lower metal stars such as Pop II and even those with $Z<10^{-4} \\ {\\rm Z}_\\odot$, but not completely metal-free stars, Pop III.1 and Pop III.2 stars. The star formation rate expressed by Eq.~(1) in \\cite{Dominik2013} is completely different from the one shown in Fig.~8 of the present paper. In the latter case, the star formation rate at $z=0$ is zero, while in the former case, it is the present star formation rate of our Galaxy which is not zero.  In Fig.~6 of \\cite{Dominik2013},  they show the cumulative merger rate as a function of redshift $z$ for different four models which corresponds to Fig.~10 of the present paper.  In our Fig.~10, for the second~(third) generation gravitational wave detectors,  $ z\\sim 0.3\\ (3)$ is the detection range, respectively. In Fig.~6 of \\cite{Dominik2013}, information on the detectability is not available, since the chirp mass distribution function is not available. Assuming  it is  similar to that in \\cite{Dominik2012}, the merger rate for the second and third generation gravitational wave detectors is either higher or lower than  Fig.~10 of the present paper taking into account that the chirp mass of Pop I and Pop II BH-BHs is smaller than that for Pop III.1 and Pop III.2 BH-BHs. If the detection rate of the coalescing Pop I and Pop II BH-BHs is lower than that of Pop III, it might be possible to confirm the existence of the massive Pop III stars by detecting gravitational waves from their remnant black hole and identifying the typical chirp mass $\\sim 30\\ \\msun$ and the cumulative redshift distribution which depends almost only on the cosmological parameters. On the other hand, if the detection rate of the coalescing Pop I and Pop II BH-BHs is higher, Pop III BH-BHs contribute only some parts of the gravitational wave events of BH-BHs. If the detected number is $\\sim 10^4$ for the third generation detector like ET, we might select a Pop III BH-BH from its mass and be able to draw the cumulative redshift distribution function to confirm the existence of Pop III stars. In any case, it is needless to say that there are many undetermined parameters and distribution functions for the Pop III population synthesis so that more theoretical study on the evolution and initial conditions of Pop III binaries including the star formation rate is urgent."
    },
    "1402/1402.0959_arXiv.txt": {
        "abstract": "Neutrino emissions from electron/positron capture on the deuteron and the nucleon-nucleon fusion processes in the surface region of a supernova core are studied. These weak processes are evaluated in the standard nuclear physics approach, which consists of one-nucleon and two-nucleon-exchange currents and nuclear wave functions generated by a high precision nucleon-nucleon potential. In addition to the cross sections for these processes involving the deuteron, we present neutrino emissivities due to these processes calculated for typical profiles of core-collapsed supernovae. These novel neutrino emissivities are compared with the standard neutrino emission mechanisms. We find that the neutrino emissivity due to the electron capture on the deuteron is comparable to that on the proton in the deuteron abundant region. The implications of the new channels involving deuterons for the supernova mechanism are discussed. ",
        "introduction": "\\label{sec:introduction}  The neutrinos play pivotal roles in core-collapse supernovae and their subsequent evolution to neutron stars. Neutrino reactions in the dense matter % of a supernova core are crucial for understanding the supernova explosion mechanism, which is still elusive despite extensive studies over decades. It is therefore essential to identify all neutrino processes, both neutrino-emission and neutrino-absorption processes, that can be important in the supernova environment. Failing to include all the relevant neutrino processes in supernova modeling may have significant consequences for the theoretical understanding of the supernova explosion. The emission of neutrinos acts as a  cooling mechanism of the central core, and the addition of any neutrino emission channels that have not been considered so far could affect this cooling mechanism. A portion of the emitted neutrinos are subsequently absorbed by the material behind the shock wave and thereby act as a heating agent. Additional sources of neutrinos enhance this neutrino-heating mechanism and may help the revival of the  stalled shock wave and lead to a successful modeling of supernova explosion \\citep{bet90,kot06,jan07}. The neutrinos emitted gradually ($\\sim$20 s) from a nascent neutron star (proto-neutron star) in a supernova explosion, can be detected as supernova neutrinos at terrestrial neutrino detectors, like in the case of SN1987A \\citep{suz94}, and % can be a useful source of information about the neutrino emission mechanisms.   Recent calculations have shown that deuterons, tritons and $^{3}$He can appear copiously in the regions between the supernova core and the shockwave \\citep{sumroe08,arc08,hemp12}. These light elements have so far not been included in the tables of equation of state (EOS)~\\citep{lat91,she98a,she98b} that are routinely used in supernova simulations where the nuclear species are limited to the proton, neutron, $^{4}$He and one ``representative heavy nucleus\" that is assumed to simulate the roles of all heavy nuclei. The light elements with mass number $A\\!=\\!2\\,{\\rm and}\\,3$ can be abundant in hot and moderately dense matter ($  < 10^{13}g/cm^{3}$) under nuclear statistical equilibrium~\\citep{sumroe08,arc08,hemp12,furusawa13a} and should be considered in studying the supernova mechanism. They appear in the heating region behind a shockwave, and also in the cooling region at the surface of a proto-neutron star; their appearance gives a new contribution to the neutrino opacity. For example, ~\\citet{sumroe08} showed in a snapshot 150 msec after the bounce that  the deuteron mass fraction amounts to about 10\\%   % in the neutrino-emitting cooling regions at densities $\\sim$10$^{11}$-10$^{12}$ g/cm$^{3}$. Neutrino processes in this cooling region are essential for determining the flux and spectra of emitted neutrinos, which in turn affect the efficiency of neutrino heating behind the shockwave. Nakamura et al.~\\citep{nakamura09} investigated neutrino absorption on deuterons as an additional heating mechanism on top of the neutrino reactions on nucleons and $^{4}$He, while Arcones et al.~\\citep{arc08} studied neutrino reactions on tritons and $^{3}$He to evaluate their influences on the neutrino spectra at the outer layer of a proto-neutron star.  According to Refs.~\\citet{sumroe08,arc08}, the deuteron fraction can be larger than the proton fraction in part of the neutrino-sphere region between the shock wave and the proto-neutron star surface. This indicates that weak-interaction deuteron breakup may play a significant role in neutrino emission processes, possibly altering the conventional understanding of the role of the protons in the neutrino-emission processes as well as the neutronization of matter. An additional reaction to be investigated in this work is deuteron formation in nucleon-nucleon scattering, which also leads to neutrino emission. Although both these neutrino emission processes certainly exist on top of the conventional neutrino-emission processes, they have so far not been considered in supernova simulations.  In this article we study neutrino emissions from \\underline{d}euteron \\underline{b}reakup/\\underline{f}ormation processes (DBF for short) in the surface region of a proto-neutron star, where neutrino emissions act as a cooling mechanism. We present the first evaluation of the neutrino emissivities from electron/position-capture on the deuteron (deuteron breakup) and from the nucleon-nucleon weak fusion processes (deuteron formation); see (\\ref{eqn:eled})-(\\ref{eqn:pnd}) below. The neutrino emissivities arising from DBF will be compared with those coming from the ``conventional\" processes; by ``conventional\" processes we mean the neutrino emission processes which have been previously considered in the literature, and which are listed in (\\ref{eqn:elep})-(\\ref{eqn:epem}) below. The neutrino emissivities due to DBF reported here are expected to be useful for numerical simulations of supernova explosion and proto-neutron star cooling.  Theoretical treatments of electroweak processes in two-nucleon systems % are well developed. For low-energy neutrino-deuteron reactions, serious efforts to reduce theoretical uncertainties have been made in order to analyze data from the Sudbury Neutrino Observatory~\\citep{nsk01,net02}. One approach is the standard nuclear physics approach (SNPA) that involves nuclear wave functions derived from high-precision phenomenological nuclear potentials, and one-nucleon and two-nucleon electroweak currents. SNPA has been well tested by analyses of photo-reactions, electron scattering, and muon capture on the two-nucleon systems~\\citep{doi90,tamu92,sato95}. Another theoretical approach, effective field theory (EFT) consisting of nucleons and pions, has been developed and applied to low-energy electroweak processes~\\citep{kubo04}. Both methods essentially agree with each other for low-energy electroweak processes in the two-nucleon systems. The $pp$-fusion process, $pp\\to d e^- \\nu_e$, is one of such processes relevant to this work. This reaction has been studied with both SNPA and EFT, and good agreement between the two methods has been found~\\citep{sch98,tsp03}. Another nucleon-nucleon fusion process relevant to this work is neutron-neutron fusion, which was previously studied with EFT~\\citep{ando06}. In the present work we adopt SNPA.  This article is arranged as follows. In section \\ref{sec:reaction} we discuss neutrino emissions via DBF relevant to the supernova environment. The theoretical framework for calculating the cross sections for neutrino emission via DBF and the corresponding neutrino emissivities are outlined in section \\ref{sec:formulation}, and the numerical results are presented in sections \\ref{sec:cros} and \\ref{sec:SecEmissivity}. The implications and a summary of these results for the supernova mechanism are discussed in section~\\ref{sec:summary}.    % ",
        "conclusions": "\\label{sec:summary}  It was pointed out in Refs.~\\citet{sumroe08} and \\citet{arc08} that, in addition to deuterons, tritons can also have large abundances in high density regions (10$^{11}$--10$^{14}$ g/cm$^{3}$), where the electron fraction $Y_e$ is low and the temperature $T$ is high. This suggests the possible importance of neutrino emissivities involving the triton or ``triton\" (triton-like three-nucleon correlation in dense matter). In the present work, however, we have not considered these effects.  Our study here is based on the spherical (1D) configurations of supernovae. It would be interesting to study neutrino emission with the abundance of light elements in 2D/3D profiles. Hydrodynamical instabilities can generate non-spherical distribution of matter in the cooling region around a proto-neutron star and the heating region behind a stalled shock wave. Since the density, temperature and electron fractions can have wider ranges in 2D/3D profiles, there may be regions of high deuteron abundance that cannot be found in the 1D profile. The existence of deuterons in the heating regions can contribute to the additional source of heating as studied by \\citet{nakamura09}. It thus seems interesting to study the possible effects of the neutrino emission and absorption channels involving the deuteron in multi-D supernova explosion simulations. In principle, one must study effects of all neutrino processes by solving the neutrino transfer and hydrodynamics, with detailed information on composition from the equation of state of supernova matter. A study along this line has been recently made~\\citep{furusawa13}, taking into account the neutrino processes involving neutrino absorption on the deuteron and the light element abundance.  We now summarize. Neutrino emissions from $e^\\mp$-capture on the deuteron and from deuteron formation in nucleon-nucleon weak-fusion processes have been studied as new neutrino emission mechanisms in supernovae. These weak processes are evaluated with the standard nuclear physics approach, which consists of the one-nucleon impulse current and two-nucleon exchange current and nuclear wave functions derived from high-precision phenomenological $NN$ potentials. It is found that the contribution of the two-nucleon meson-exchange current is only a few \\% for the $e^\\mp$-capture reactions, while it can be as large as the one-nucleon current contribution for the NN fusion reaction at higher energies. The consequences of these new neutrino-emission channels have been examined for representative profiles of core-collapse supernovae at 150 ms after core bounce. The emissivity due to the $e^\\mp$ capture reaction on the deuteron is found to be smaller than that on the free nucleon. Therefore, as Fig.~\\ref{fig:ecap-ratio} indicates, the total neutrino emissivity due to electron capture on protons and deuterons is suppressed when an appreciable amount of protons in a supernova are bound inside deuterons. This results in a smaller neutrino luminosity and the lower efficiency of neutrino heating behind a stalled shock wave. Therefore, this new process may contribute unfavorably towards a successful supernova explosions. It might lead to a slower speed of the deleptonization and, hence, a slower evolution of nascent proto-neutron stars. On the other hand, as seen in Figs.~\\ref{fig:ecap-out} and \\ref{fig:nn-out}, neutrino emission via deuteron formation can be comparable to nucleon-nucleon bremsstrahlung in the outer region. This implies that there might exist situations in which the deuteron-formation weak processes are the main channels for neutrino emission.  In the inner core region, where the electrons are highly degenerate (high densities at low temperatures), pair-production via $e^-e^+$ annihilation is suppressed, making nucleon-nucleon bremsstrahlung a main channel to produce $\\nu\\bar{\\nu}$ pairs among the conventional processes\\citep{suz93, bur06}.  % Meanwhile,  ``deuteron\" formation processes in $NN$ scattering can have large rates for $\\nu_\\mu$ and $\\nu_\\tau$ emissions, a feature that may have significant consequences for the cooling of compact stars. Furthermore, the possible modification of the energy spectra of $\\nu_e$, $\\nu_\\mu$ and $\\nu_\\tau$ due to ``deuteron\" formation may influence supernova nucleosynthesis~\\citep{woo90,yoshi04} and the terrestrial observation of supernova neutrinos \\citep{nakazato13}. On the other hand, the possible increase of the $\\nu_\\mu$ and $\\nu_\\tau$ fluxes due to ``deuteron\" formation hardly affects the heating process behind a shockwave, because these low-energy $\\nu_\\mu$ and $\\nu_\\tau$ interact with stellar matter only through the NC. As explained earlier, the ``deuteron\" here stands for a tensor-correlated $NN$ pair that may persist even in dense nuclear matter. A detailed study of deuteron-like two-nucleon tensor correlation in dense matter seems well warranted, but it is beyond the scope of our present exploratory work."
    },
    "1402/1402.4057_arXiv.txt": {
        "abstract": "{We describe the late spectral variability and flux evolution of TDF Sw J1644+57, a Tidal Disruption Flare which left the typical potential trend proportional to $t^{-5/3}$ in 2012, maintaining a quiescent flux until nowadays. Sixteen X-ray observations of ESA satellite \\emph{XMM-Newton} have been used in this study, including the one performed on 17$^{\\rm th}$July, 2013. A search for optical emission in BOOTES/CASANDRA database has been performed too. Late X-ray fluxes show that the source flux decline does not follow the expected TDF trend at the time of the last \\emph{XMM-Newton} observation. Moreover, the spectra fitting parameters, in particular the neutral hydrogen column density, N$_{\\rm H}$, and the power-law index, $\\Gamma$, indicate that the source darkening has diminished and that the spectral shape has flattened with time. The disruption of the star could have come to an end. Nevertheless, a quiescent X-ray flux continues. Evidence for a quiescent X-ray flux is presented.}  \\resumen{Presentamos la variabilidad espectral tard\\'ia y la evoluci\\'on de flujo del TDF Sw J1644+57, un fen\\'omeno transitorio probablemente producido por el desgarramiento de una estrella por fuerzas de marea, el cual en 2012 abandon\\'o la tendencia caracter\\'istica proporcional a $t^{-5/3}$ para mantener un flujo estable en quietud hasta la actualidad. Para este trabajo se han empleado diecis\\'eis observaciones del sat\\'elite de la ESA \\emph{XMM-Newton}, incluida la \\'ultima observaci\\'on (17 de Julio de 2013). Asimismo, se ha llevado a cabo una b\\'usqueda en la base de datos de BOOTES/CASANDRA a fin de detectar la contrapartida \\'optica. Los flujos tard\\'ios en rayos X muestran que el decaimiento de flujo de la fuente ya no sigue la tendencia esperada para un TDF. Adem\\'as, los par\\'ametros del ajuste espectral, en particular la densidad de columna de hidr\\'ogeno neutro, N$_{\\rm H}$, y el \\'indice de la ley de potencias, $\\Gamma$, indican que el oscurecimiento de la fuente ha disminuido y que los espectros se han aplanado con el tiempo. El desgarramiento de la estrella podr\\'ia haber llegado a su fin. Sin embargo, el flujo de rayos X en quietud contin\\'ua. Hay indicacaciones de emisi\\'on en quietud.}   \\addkeyword{Galaxies: Active} \\addkeyword{ISM: Jets and outflows} \\addkeyword{Stars: Flare} \\addkeyword{X-rays: Galaxies}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "\\label{sec:Results}   \\begin{asparaitem}  \\item The  N$_{\\rm H}$ has decreased as  Figure~\\ref{fig:NH} shows, between day 188 and 549, which indicates that most of the clouds of gas and dust around the disrupted star have dissipated. The final number of 0.056$\\times10^{22}$cm$^{-2}$ is very close to the galacticN$_{\\rm H}$, and implies that the phenomenon is less energetic now.  \\item The spectral shapes have flattened with time, as Figure~\\ref{fig:Flattening} reveals when we compare the spectrum of the first observation, performed on day 3, with the last one, taken on day 842. Besides, in late observations the maximum flux occurs at a lower energy.  \\begin{figure}[!t] \\includegraphics[width=\\columnwidth]{Intrinsic_Flux} \\caption{Evolution of the intrinsic flux (0.5$-$8.0 keV). \\emph{Swift} observations are drawn as black circles. The most recent \\emph{Chandra} observation is shown as a black square. The fluxes obtained in this work are presented as orange circles. In red, at the bottom right, our latest \\emph{XMM-Newton} observation. Adapted from \\citet{2013ApJ...767..152Z}} \\label{fig:Intrinsic_Flux} \\end{figure}  \\item The intrinsic fluxes between 0.5 and 8.0 keV calculated from the sixteen observations of the \\emph{XMM-Newton} Archive, are consistent with previous works about Sw J1644+57. See  Figure~\\ref{fig:Intrinsic_Flux}.  \\item Between late 2011 and 2012, the source X-ray flux started to decline, leaving the characteristic potential tendency of a TDF in 2012. Consequently, the black hole has probably used up the star that provided dust and gas clouds, or its remains have escaped from the orbit around the hole \\citep{1988Natur.333..523R}.  \\item The object behavior has changed from an X-ray emission compatible with a TDF, to a quiescent emission that remains until nowadays and that is higher than Sagittarius A$^{*}$. We ignore if the accretion already existed before \\emph{Swift} triggered.  Thus, the source could be the beginning of an Active Galactic Nucleus, AGN, with the jet faced on, i.e., a mini-blazar. Additional X-ray observations are needed to confirm it.  \\item We do find no transient optical emission (brighter than 10$^{\\rm th}$ mag) at the CASANDRA all-sky images taken at the different BOOTES robotic astronomical stations world wide in the period 2009-2013 when data is available \\citep{2012ASInC...7..313C}.  \\end{asparaitem}"
    },
    "1402/1402.1444_arXiv.txt": {
        "abstract": "We formulate a model of pulsar spin evolution (braking, inclination angle evolution and radiative precession) taking into account the non-rigidity of neutron star rotation. We discuss two simple limiting cases of this model and show that the evolution of the inclination angle substantially depends on the model of crust-core interaction. The non-rigidity of core rotation accelerates the inclination angle evolution and makes all pulsars evolve to the orthogonal state. The size of the effect depends on the amount of differentially rotating matter and mechanism of its interaction with the rest of the star. Since the rapid evolution of the inclination angle apparently contradicts the observational data, our results may be used as an additional test for the theories of the cores of neutron stars. ",
        "introduction": "A rotating magnetized neutron star, if it has perfectly spherical shape, can be characterized by two vectors: angular velocity vector $\\vec{\\Omega}$ and magnetic moment vector $\\vec{m}$. During the life of a neutron star the magnitudes of these vectors as well as the inclination angle $\\chi$ formed by these vectors evolve in time. This evolution is caused by electromagnetic torque acting on rotating magnetized star.  A neutron star is surrounded by a large magnetosphere. Strictly speaking, the electromagnetic torque can be calculated only by using a self-consistent theory of this magnetosphere which despite the achieved progress is far from complete at present \\cite{Spitkovsky2008}.   The problem becomes even more complicated if one wishes to take into account the internal structure of neutron star. Neutron stars are not perfectly rigid. They contain a liquid core. There are two main mechanisms of crust-core interaction: coupling through the magnetic field and viscosity \\cite{Easson1979}. The first mechanism is much more effective and, if it works, then the protons, electrons and normal neutrons in the core can be considered  as rigidly rotating with the crust. In contrast, the superfluid neutrons which are believed to be present in neutron star's core \\cite{YakovlevLevenfishShibanov1999} are decoupled from the rest of the matter and interact with it only by weak mutual friction force (vortex-mediated interaction) \\cite{HallVinen1956}.  The configurations when the magnetic field does not penetrate the core have been discussed in the literature \\cite{PonsMirallesGeppert2009,GourgouliatosEtAl2013}. The expulsion of magnetic field can be caused, for example, by type-I superconductivity of core protons. Despite the fact that the type-II proton superconductivity is more likely from the point of view of microscopic calculations (see, however, \\cite{BuckleyMetlitskiZhitnitsky2004a}), its implications for the  neutron star rotation dynamics are not so clear. The coexistence of proton magnetic fluxoids with neutron superfluidity seems to be inconsistent with the observed long periods of precession in some neutron stars \\cite{Link2006}.  In the present paper we consider in detail the two limiting cases of strong and weak coupling between the crust and the core of a neutron star, which correspond to the two physical cases discussed above. ",
        "conclusions": "One can see that the evolution of the inclination angle dramatically differs for different crust-core interaction models. The rate of the evolution depends on the amount of the matter which rotates differentially. Moreover, all trajectories sharply turn upwards when the action of anomalous torque becomes comparable with the action of normal torque. These features should be included in the studies of population synthesis of pulsars. They can also be used as additional tests for the theories of neutron star interiors."
    },
    "1402/1402.3394_arXiv.txt": {
        "abstract": "Using data from the SDSS-DR7, including structural measurements from 2D surface brightness fits with GIM2D, we show how the fraction of quiescent galaxies depends on galaxy stellar mass $M_*$, effective radius $R_e$, fraction of $r-$band light in the bulge, $B/T$, and their status as a central or satellite galaxy at $0.01<z<0.2$.  For central galaxies we confirm that the quiescent fraction depends not only on stellar mass, but also on $R_e$.  The dependence is particularly strong as a function of $M_*/R_e^\\alpha$, with $\\alpha\\sim 1.5$.  This appears to be driven by a simple dependence on $B/T$ over the mass range $9<\\log(M_*/M_\\Sun)<11.5$, and is qualitatively similar even if galaxies with $B/T>0.3$ are excluded. For satellite galaxies, the quiescent fraction is always larger than that of central galaxies, for any combination of $M_*$, $R_e$ and $B/T$.  The quenching efficiency is not constant, but reaches a maximum of $\\sim 0.7$ for galaxies with $9<\\log(M_*/M_\\Sun)<9.5$ and $R_e<1$ kpc.  This is the same region of parameter space in which the satellite fraction itself reaches its maximum value, suggesting that the transformation from an active central galaxy to a quiescent satellite is associated with a reduction in $R_e$ due to an increase in dominance of a bulge component. ",
        "introduction": "It is now well known that galaxies in the local Universe can generally be separated into two broad classifications: blue, star-forming, disk galaxies and red, elliptical and lenticular galaxies with little or no recent star formation \\citep[e.g.][]{Strateva,Kauffmann3,Baldry,Balogh2}.  The star formation rate (SFR) of blue galaxies is well correlated with their stellar mass (M$_*$), forming the ``main sequence'' or ``star-forming sequence'' in the SFR-M$_*$ plane \\citep{Brinchmann,Salim,Gilbank,MBW+}.    Surveys show that the star-forming sequence has persisted out to z $=$ 1 \\citep{Bell,Willmer,Noeske} and possibly z $=$ 2 and further \\citep{Brammer,Taylor2}.  Most tellingly, the number density of red, passive galaxies has roughly doubled since z $=$ 1, while the number of blue galaxies has remained nearly constant \\citep{Bell,Faber,Brown,Goncalves,Muzzin}.  This motivates the search for processes that cause blue, star-forming galaxies to transition to red, non-star-forming ones. This transition has become known as ``quenching''.  It has also been established that the galaxy population varies with environment, in the sense that low-density environments are dominated by less massive, blue star-forming galaxies while denser groups and clusters tend to host massive red elliptical and lenticular galaxies \\citep[e.g.][]{Oemler,DavisGeller,Dressler,Baldry2,Pasquali+10}. Interestingly, the star-forming sequence itself appears to be largely independent of environment, while the fraction of star-forming galaxies decreases dramatically with increasing local density \\citep[e.g.][]{Carter,Balogh2,Wolf,P+08,Peng10}. \\begin{table*} \\begin{tabular}{cccccccccccccc} \\hline VAGC& Yang & Simard & RA & Dec & z&R$_e$&B/T&$M_*$&SFR& $m_r$&Central&C&V$_{\\rm max}$\\\\ ID & ID & Rec. no. & \\multicolumn{2}{c}{(J2000)} && (kpc)&  &($10^{10} M_\\Sun$)& ($M_\\Sun/$yr)&(mag)&flag&&($10^{6}$Mpc$^3$)\\\\ \\hline 234667 & 1129345 & 432305 & 161.86430 & 13.889750 & 0.010 & 1.00 & 0.23  & 0.02 & 0.04 & 16.34 & 1 & 0.93 & 0.315\\\\ 204290 & 958798 & 1095350 & 173.42712 & 50.449570 & 0.010 & 1.47 & 0.36 & 0.05 & 0.06 & 15.73 & 1 & 0.88 & 0.761 \\\\ 114459 & 431764 & 897676 & 151.88120 & 57.729069 & 0.010 & 1.83 & 0.01 &  0.006 & 0.04 & 16.95 & 1 & 0.89 & 0.098 \\\\ 256139 & 1219961 & 975660 & 178.49693 & 12.198710 & 0.010 & 1.87 & 0.17 &  0.08 & 0.15 & 15.08 & 1 & 0.86 & 1.875 \\\\ \\hline \\end{tabular} \\caption{A sample of entries in our catalogue, available in the online version. These data are simply matched entries from the published catalogues of \\citet{Blanton2005}, \\citet{Brinchmann}, \\citet{Simard}, and \\citet{YangDR7}.  Columns 1-3 provide the ID for simple matching with \\citet{Blanton2005}, \\citet{YangDR7} and \\citet{Simard}, respectively.  All catalogues were matched on RA, Dec and redshift (columns 4-6).  Column 7 is the circularized half-light radius from the \\citet{Simard} bulge/disk decomposition model, with free index $n$ for the bulge component; column 8 is the corresponding fraction of light in the bulge.  Columns 9 and 10 are the stellar mass and SFR from \\citet{Brinchmann}, and column 11 is the $r$-band petrosian magnitude from the original SDSS catalogue.  Column 12 is a flag that indicates whether (1) or not (2) a galaxy is the most luminous in its halo according to the \\citet{YangDR7} catalogue; galaxies with a flag of 1 are considered a ``central'' in our analysis.  Finally, column 12 is the spatial completeness from \\citet{Blanton2005}, and column 13 is the selection volume that we compute, based on the absolute magnitude and $R_e$ of the galaxy.  Additional columns that may be useful, including some that we make use of in the Appendix, are provided in the online table.\\label{tab-catalogue} }. \\end{table*}  There is some evidence to show that galaxy properties depend on the mass of their host, dark matter halo, with distinctly different behaviour for the most massive galaxies in the halo (known as centrals, since they are generally expected to lie near the centre of the halo) and the smaller, satellite galaxies that orbit around them  \\citep{Hogg,Berlind,Pasquali+10}.  Some \\citep{vdB,Peng12,Woo} suggest that the correlations with environment are due to transformations that take place when a galaxy first becomes a satellite.  This is natural, as such galaxies are expected to experience diminished cosmological gas accretion \\citep{Balogh1,Kawata}, ram-pressure with the intrahalo gas \\citep{Gunn,Grebel}, and a different type of interaction with neighbouring galaxies \\citep{Moore1,Moore2} and the tidal field \\citep{Read}.  The fraction of quiescent galaxies clearly correlates with galaxy stellar mass and this, combined with the apparent ease with which stellar masses can be computed from photometric data, has motivated modelers and observers to consider stellar mass as the fundamental parameter that determines a galaxy's properties \\citep[e.g.][]{Peng10,Woo}.  However, various authors \\citep[e.g.][]{Brinchmann,Kauffmann2,Franx} found that galaxy surface density is more predictive of the specific SFR (sSFR) than stellar mass.  Others have shown the correlation is best with velocity dispersion, either as measured directly from dynamics \\citep[e.g.][]{SLH,Wake2,Wake1}, or inferred from the virial relations between mass and size \\citep{Franx}. There are good indications that the trends are driven by conditions in the centre of galaxies, either as inferred from  S\\'{e}rsic indices derived from single-profile fits to the surface brightness \\citep{Blanton2003c,Driver,Allen,Schiminovich,Bell2,Donofrio+11}, central surface density \\citep[e.g.][]{Cheung,Fang}, or presence of a central bulge \\citep{Martig,Cappellari2}. While it remains controversial which of these characterisations is most fundamental, it seems clear that galaxy structure, and not just stellar mass, is relevant \\citep[e.g.][]{P+13}.  If galaxy properties are primarily determined by some aspect of their structure, rather than their stellar mass alone, it reopens the question of how (or if) galaxies of a given {\\it structure} depend on their environment.  The first purpose of this paper is to quantify the quenching efficiency of satellite galaxies as a function of their structure.  This leads us to take a more general approach than some of the studies listed above, and consider how the quiescent fraction of central galaxies depends on galaxy stellar mass and radius in the local Universe sampled by the Sloan Digital Sky Survey \\citep[SDSS,][]{SDSS}.  All calculations use the $\\Lambda$CDM cosmology consistent with the nine year data release from the WMAP mission, with parameters $h$ = 0.693, $\\Omega_M=0.286$, and $\\Omega_\\Lambda=0.714$ \\citep{WMAP}. ",
        "conclusions": "We have used a sample of 471,554 galaxies at 0.01 $<$ z $<$ 0.2 selected from the SDSS DR7, to examine how the fraction of quiescent galaxies depends upon stellar mass, effective radius, bulge fraction, and whether a galaxy is central or satellite in the group catalogue of \\citet{YangDR7}.  We use the GIM2D parametric fits of \\citet{Simard} to determine the circular half-light radius ($R_e$) of all galaxies in the sample, and the specific star formation rates and stellar masses of \\citet{Brinchmann} to classify galaxies as active (on the star--forming main sequence) or quiescent.  We make the following observations:  \\begin{enumerate}  \\item The fraction of quiescent galaxies is a steeper function of M$_*$/R$_e$ and M$_*$/R$_e^2$ than it is of stellar mass alone.  This confirms the findings of many other authors \\citep[e.g.][]{Kauffmann2,Franx,Wake2,Wake1,Cheung}.  \\item The quiescent fraction of satellite galaxies is always higher than that of central galaxies, except perhaps at the highest masses where almost all galaxies are quenched centrals.  Thus a galaxy's status as satellite or central in its halo plays a role in its evolution.  The quenching efficiency \\citep{Peng10}, which attempts to quantify the fraction of satellite galaxies that had their star formation quenched as a satellite, is $\\sim 40$ per cent, approximately independent of stellar mass for $M_*>10^9M_\\Sun$.  \\item We presented the quiescent fraction of central galaxies as function of both $M_\\ast$ and $R_e$ in Figure~\\ref{fig:qfd}.  This representation makes it clear that the fraction does not depend on stellar mass alone, and that over most of the parameter space it appears to depend most strongly on $\\sim$M$_*$/R$_e^{1.5}$.  \\item We showed that the quiescent fraction of central galaxies correlates strongly with $B/T$ in a simple way that largely accounts for the more complex dependence on $R_e$. Galaxies with more than 40 per cent of their light in a bulge component are almost all quiescent, and galaxies less massive than $M_*=10^9M_\\Sun$ are almost all forming stars.  Otherwise, the quiescent fraction increases monotonically with increasing B/T at fixed $M_*$.  \\item The quenching efficiency is highest for galaxies with 10$^{9}<$ M$_*<$ 10$^{10}$ M$_\\Sun$ and 10$^{-0.2}<$ R$_e<$ 10$^{0.2}$ kpc (Figure~\\ref{fig:qeff}), but appears roughly constant above a certain threshold when taken solely as a function of M$_*$, M$_*$/R$_e$, or M$_*$/R$_e^2$.  \\item There are structural differences between satellite and central galaxies, such that, at a given stellar mass, satellite galaxies tend to have smaller $R_e$ and more light in a bulge component.  The fraction of satellites is greatest (at about 50 per cent) for galaxies with $9.0<\\log({M_*}/M_\\Sun)<10.0$, and $R_e<2$kpc.    Even after accounting for this difference by matching central and satellite galaxies to have the same $M_*$, $R_e$, and B/T, the quenching efficiency remains positive at all masses, and reaches a maximum of 70 per cent for galaxies with $\\log(M_*/M_\\Sun)\\sim 9.4$ and $R_e\\sim$1 kpc.  \\end{enumerate}  For central galaxies, it is appealing to consider a probability of star formation activity ending that depends on their internal structure.  In particular the strong correlation between quiescent fraction and  $B/T$ points to the dominance of a central bulge component as the driving parameter. We showed that, amongst star--forming central galaxies with $M_*<10^{11}M_\\Sun$, the average sSFR decreases with increasing $M_*/R_e$, qualitatively suporting a picture in which the presence of a bulge  reduces the sSFR and increases the probability of being quenched.  Alternatively, the data could be explained if the probability of quenching depends only on $M_*$, but is accompanied by a change in structure, with $R_e$ shrinking by a factor $\\sim 2$ due to the increased prominence of a bulge when star formation is terminated.  For satellite galaxies, that fact that the quenching efficiency is a maximum for the same combination of $M_*$ and $R_e$ at which the satellite fraction itself peaks almost certainly means that their $R_e$, and perhaps their mass, changes upon quenching."
    },
    "1402/1402.4782_arXiv.txt": {
        "abstract": "{ The knowledge of atmospheric parameters --~such as temperature, pressure, and humidity~-- is very important for a proper reconstruction of air showers, especially with the fluorescence technique. The Global Data Assimilation System (GDAS) provides altitude-dependent profiles of these state variables of the atmosphere and several more. Every three hours, a new data set on 23~constant pressure level plus an additional surface values is available for the entire globe. These GDAS data are now used in the standard air shower reconstruction of the \\pao. The validity of the data was verified by comparisons with monthly models that were averaged from on-site meteorological radio soundings and weather station measurements obtained at the Observatory in \\mal. Comparisons of reconstructions using the GDAS data and the monthly models are also presented. Since GDAS is a global model, the data can potentially be used for other cosmic and gamma ray detectors. Several studies were already performed or are underway for several locations worldwide. As an example, a study performed in Colorado as part of an Atmospheric \\rnd for a possible future cosmic ray observatory is presented. } ",
        "introduction": "}  A cosmic ray particle entering the atmosphere can initiate an extensive air shower. The secondary shower particles excite nitrogen molecules in the air which emit a characteristic, isotropic emission in the UV range as part of their de-excitation process. The light can then be observed by an optical telescope, typically consisting of a collecting mirror and a camera. To properly reconstruct the properties of such air showers, the atmospheric conditions at the site have to be known in order to correct for Rayleigh scattering effects and to estimate the fluorescence yield of the air shower~\\cite{Abraham:2010atmo}. Height-dependent profiles of temperature, pressure and humidity as well as weather conditions near the ground are relevant.  The \\pao~\\cite{Abraham:2004auger} is a cosmic ray detector located near \\mal in the Mendoza province in Argentina. It consists of a Surface Detector (SD) array and five Fluorescence Detector (FD) buildings~\\cite{Abraham:2009fd}. Between 2002 and 2010, atmospheric conditions over the Observatory were measured by intermittent meteorological radio soundings. Additionally, ground-based weather stations measure surface data continuously in order to provide the atmospheric parameters to properly reconstruct the measured air showers.  In south-east Colorado, several balloon soundings were performed as part of an atmospheric \\rnd project. The aim of this effort was to study possible enhancements and performance improvements for the \\pao, as well as explore technological advancements for a possible future ground-based observatory. The ground station used for the soundings was a mobile and slightly advanced version of the equipment used in Argentina. The launches were performed at two sites, the Atmospheric Monitoring Telescope (AMT) and the Distant Raman Laser Facility (DRLF)~\\cite{Wiencke:2011icrc}. The sites are about 40\\,km apart and are both equipped with identical weather stations.  Performing radio soundings imposes a large burden, both in terms of funds and manpower. We investigated the possibility of using data from the Global Data Assimilation System (GDAS)~\\cite{GDASinformation}, a global atmospheric model, for the site of the \\pao~\\cite{Abreu:2012gdas,Will:2011icrc}.  GDAS data are publicly available free of charge via READY (Real-time Environmental Applications and Display sYstem). Each data set contains all the main state variables as a function of altitude. The data gathered in Colorado were also compared to GDAS data in order to evaluate the possibility to use GDAS also in different locations and for a possible future ground-based cosmic ray detector. ",
        "conclusions": "}  The comparison of GDAS data for the site of the \\pao in Argentina with local atmospheric measurements validated the adequate accuracy of the 3-hourly GDAS data. An air shower reconstruction analysis confirmed the applicability of GDAS for Auger reconstructions and simulations, giving improved accuracy when incorporating GDAS data instead of MM. Also, the value of using an atmosphere-dependent fluorescence description has been demonstrated. For the Colorado \\rnd site, the differences between the measured radiosonde data and GDAS are of the same order as in Argentina, further supporting the general validity of GDAS data as an atmospheric description to be used in current and future cosmic ray observatories.  \\vspace*{0.5cm} \\footnotesize{{\\bf Acknowledgment:} We would like to thank the organizers of the workshop \\emph{AtmoHEAD: Atmospheric Monitoring for High-Energy Astroparticle Detectors} in Saclay, France, 2013 for the inspiring meeting. Part of these investigations are supported by the Bundesministerium f\\\"ur Bildung und Forschung (BMBF) under contracts 05A08VK1 and 05A11VK1. Furthermore, these studies would not have been possible without the entire Pierre Auger Collaboration and the local staff of the \\pao.{}}"
    },
    "1402/1402.1391_arXiv.txt": {
        "abstract": "The oscillations in solar like stars are described in terms of  the phase shifts of the eigenmodes from simple sine-waves. We discuss model fitting and inversion techniques based on this representation. We analyse the periodic signatures from the HeII ionisation zone and base of the convective envelope of the CoRoT star  HD49933. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1672_arXiv.txt": {
        "abstract": "We present an algorithm for the simultaneous measurement of a pulse time-of-arrival (TOA) and dispersion measure (DM) from folded wideband pulsar data.  We extend the prescription from Taylor (1992) to accommodate a general two-dimensional template ``portrait'', the alignment of which can be used to measure a pulse phase and DM.  We show that there is a dedispersion reference frequency that removes the covariance between these two quantities, and note that the recovered pulse profile scaling amplitudes can provide useful information.  We experiment with pulse modeling by using a Gaussian-component scheme that allows for independent component evolution with frequency, a ``fiducial component'', and the inclusion of scattering.  We showcase the algorithm using our publicly available code on three years of wideband data from the bright millisecond pulsar J1824-2452A (M28A) from the Green Bank Telescope, and a suite of Monte Carlo analyses validates the algorithm.  By using a simple model portrait of M28A we obtain DM trends comparable to those measured by more standard methods, with improved TOA and DM precisions by factors of a few.  Measurements from our algorithm will yield precisions at least as good as those from traditional techniques, but is prone to fewer systematic effects and is without ad hoc parameters.  A broad application of this new method for dispersion measure tracking with modern large-bandwidth observing systems should improve the timing residuals for pulsar timing array experiments, like the North American Nanohertz Observatory for Gravitational Waves. ",
        "introduction": "\\label{intro}  The practice of pulsar timing attempts to model the rotation of a neutron star by phase-connecting periodic observations of its pulsed, broadband radio signal.  The earliest demonstration of long-term timing observations came relatively soon after the discovery of pulsars~\\citep{Roberts&Richards71, Hewish68}.  The scientific merits garnered from pulsar timing span astrophysical fields such as planetary science, the interstellar medium, nuclear physics, gravitational wave physics, and are all well-documented (for a review, see eg. Chapter 2,~\\citet{L&K04}).  Pulsar timing and its related experiments have carved out a ``sweet spot'' in the radio frequency regime that naturally emerges as a trade-off between the pulsars' steep power-law spectra at the high-frequency end, and the low-frequency drawbacks arising from the pulsed radio signal having to propagate through the ionized interstellar medium (ISM) and the Earth's ionosphere, as well as having to compete with the diffuse background of the galactic synchrotron continuum.  The latter has a spectral index in the 1--10~GHz range of $\\approx-2.8$~\\citep{Platania98}.  Population studies have shown that pulsars have an average spectral index around -1.4 at gigahertz frequencies~\\citep{Bates13}.  The most relevant ISM effect arises from propagation through a homogeneously ionized medium.  Interstellar dispersion alters the group-velocity of the radio signal, retarding the arrival of pulses emitted at a radio frequency $\\nu$ by a time $t_{\\textrm{DM}}$ (relative to an infinite frequency signal) according to the cold-plasma dispersion law, \\begin{equation} \\label{DMlaw} t_{\\textrm{DM}} = K \\times \\textrm{DM} \\times \\nu^{-2}, \\end{equation} where $K \\equiv \\frac{e^2}{2\\pi m_ec} = 4.148808(3) \\times 10^{3}$~MHz$^2$~cm$^{3}$~pc$^{-1}$~sec is called the dispersion constant\\footnote{$K$ is a combination of the electron charge $e$, electron mass $m_e$, and speed of light $c$.  It is common practice in the pulsar community to adopt the approximation $K^{-1} = 2.41\\times10^{-4} $~MHz$^{-2}$~cm$^{-3}$~pc~sec$^{-1}$~\\citep{L&K04}, which we have used in~\\S\\ref{m28a} and~\\S\\ref{mcs}.}, and DM is the dispersion measure.  The dispersion measure is defined as \\begin{equation} \\label{DM} \\textrm{DM} \\equiv \\int\\limits_l n_e\\,dl, \\end{equation} which is the free-electron column density along the path-of-propagation $l$ to the pulsar.  The pulse-broadening effect of multi-path propagation through a turbulent, inhomogeneous ISM, known as interstellar scattering, has an even stronger spectral index $\\approx-4$, and becomes increasingly important at lower frequencies for the highest-DM, farthest pulsars~\\citep{L&K04}.  Scattering not only broadens the pulsed signal, but delays an intrinsically sharp pulse by an amount roughly proportional to its width, and so is a source of bias in timing measurements.  The determination of dispersion measures and effects from scattering have been non-trivial problems concomitant with timing measurements since the beginning~\\citep{Rankin70, Rankin71}.  Nearly all observations taken for (high-precision) pulsar timing experiments are taken within the radio window mentioned above, which lies somewhere in the two decades bounded by about 100~MHz and 10~GHz.  The middle decade centered around 1500~MHz seems to be the perennial favorite for timing experiments.  Recent developments in pulsar instrumentation and computing over the last 5--10 years have enabled more accurate and sensitive timing measurements.  Namely, coherent dedispersion, which completely removes the quadratic time-delay due to a known amount of interstellar dispersion~\\citep{Hank&Rick75}, required significant advances in computer technology before becoming feasible in real-time on a wide-bandwidth signal.  Historically, observations that implemented coherent dedispersion were limited by computing resources to a bandwidth of order $\\sim$100~MHz or less, which is less than most receiver bandwidths.  Thus, if one wanted to cover a large portion of the pulsar spectrum, either for timing, spectral, or interstellar medium purposes, several adjacent receiver bands had to be observed separately, which often meant asynchronous measurements and non-contiguous frequency coverage.   The implementation of real-time coherent dedispersion to large, instantaneously observed bandwidths has led to the regime wherein the receiver bandwidth (BW) is a limiting factor.  The first generation of GHz-bandwidth, coherent dedispersion instruments has been proliferating in the pulsar community for the past several years, beginning with the Green Bank Ultimate Pulsar Processing Instrument (GUPPI)\\footnote{\\url{www.safe.nrao.edu/wiki/bin/view/CICADA/NGNPP}} outfitted for the 100-m Robert C. Byrd Green Bank Telescope (GBT)\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}~\\citep{Duplain08}.  GUPPI is an FPGA- and GPU-based system capable of real-time coherent dedispersion of an 800~MHz bandwidth.  The smearing $\\delta t_{\\textrm{DM}}$ incurred from incorrectly dedispersing a narrow frequency-channel of bandwidth $\\Delta\\nu = \\frac{\\textrm{BW}}{n_{chan}}$ and center frequency $\\nu_c$ by an amount $\\delta{\\textrm{DM}}$ goes as \\begin{equation} \\begin{split} \\label{DMsmear} &\\delta t_{\\textrm{DM}} \\approx \\frac{2K\\ \\delta{\\textrm{DM}}\\ \\Delta\\nu}{\\nu_c^3} \\\\ &\\qquad \\approx 8.3\\ \\Big(\\frac{\\delta{\\textrm{DM}}}{\\textrm{cm}^{-3}\\ \\textrm{pc}}\\Big)\\Big(\\frac{\\Delta\\nu}{1\\ \\textrm{MHz}}\\Big)\\Big(\\frac{\\nu_c}{1\\ \\textrm{GHz}}\\Big)^{-3}\\ \\mu\\textrm{s}. \\end{split} \\end{equation} This equation demonstrates why it was difficult to obtain precise, broadband measurements from millisecond pulsars (MSPs) prior to coherent dedispersion.  Incoherent dedispersion shifts individual frequency channels of the data based on the assumed DM without compensating for dispersion within each channel.  This means that $\\delta{\\textrm{DM}}$ is equivalent to the full, true DM and so the dispersive smearing $\\delta t_{\\textrm{DM}}$ can easily be a significant fraction of the pulsar's spin period ($P_{s} \\lesssim 10$~ms for MSPs).  To mitigate this problem, a large number of filterbank channels is needed, but it comes at the expense of a degraded time resolution due to Nyquist sampling constraints.  In turn, poorer time resolution means a poorer timing precision.  Moreover, tracking the dispersion measure changes in MSP observations is necessary for minimizing the timing residuals used in gravitational wave searches with a pulsar timing array (PTA)~\\citep{You07}.  Equation~\\ref{DMlaw} says that an incorrect DM of $10^{-3}$~cm$^{-3}$~pc at 1500~MHz introduces a delay of $\\sim2 \\mu$s relative to infinite frequency, which is well above the timing quality desired by PTA experiments.  As part of the Parkes Pulsar Timing Array\\footnote{\\url{www.atnf.csiro.au/research/pulsar/array/}} project~\\citep{Hobbs13}, \\citet{Keith13} developed a method to correct for inaccurate dispersion measures based on modeling the multi-frequency timing residuals.  However, the authors also postulate that more accurate DM variations could be measured from wideband receivers, which ameliorate the difficulties of aligning pulsar data taken with different receivers in different epochs.  The desire for very broadband pulsar observations (i.e. with significantly high fractional bandwidths, $\\gtrsim$1) necessitates new, unique receiver designs that can cover much of the frequency range once concatenated from disjoint observations.  Wideband receivers and their complimentary, real-time coherent dedispersion backends will quickly facilitate developments in all realms of pulsar astrophysics, including studies of the pulsar spectrum, magnetosphere, and ISM properties.  One such instrument, called the Ultra-Broad-Band (UBB) receiver and associated backend\\footnote{\\url{www3.mpifr-bonn.mpg.de/staff/pfreire/BEACON.html}}, has been recently installed at the Effelsberg 100-m Telescope and covers a frequency range from $\\sim$600~--~$\\sim$3000~MHz.  However, the current method for making pulse time-of-arrival (TOA) measurements that is used almost ubiquitously in the pulsar timing community does not use all of the information contained in new broadband observations.  In summary, the protocol employs frequency-averaged pulse profiles as models of the pulsar's signal for entire receiver bands, which ignores any profile evolution intrinsic to the pulsar or imposed by the ISM.  Both intrinsic profile evolution and DM changes are usually taken into account in the timing model for the pulsar's rotation, but there is no modeling of the effects from scattering or scintillation.  Arbitrary phase-offsets (known as ``JUMPs'') are introduced to align disparate template profiles that are used to measure TOAs from different frequency bands.  Multi-channel TOAs are also parameterized by both a quadratic delay (proportional to the DM) and an arbitrary function to remove residual frequency structure from otherwise unmodeled profile evolution.  Additionally, many multi-channel TOAs that are adversely affected by scintillation are often cut from being included in the timing model fit due to their inaccurately determined uncertainties; in effect, one is throwing away portions of the band that do contain much signal.  These methods are ad hoc and incomplete in that they were developed as the availability of bandwidth and multi-frequency observations became a ``problem'' (cf. ``the wide-bandwidth problem''~\\citep{L&D13}), and were appropriate when observations covered a narrow bandwidth: phase JUMPs account for profile evolution occurring in frequency gaps that are not observed.  It seems natural in the era of wideband receivers --- when frequency evolution {\\it is} observed in the band --- to devise a method for TOA measurement that includes a frequency-dependent model of the average pulse profile.  In doing so, it becomes straightforward to include a simultaneous measurement of the dispersion measure.  As we will show, a very simple extension to the algorithm that is currently used is a first step in a more comprehensive and necessary description of the received pulsar signal. ",
        "conclusions": "\\label{conc} In this paper, we have presented a simple method for measuring TOAs in folded pulsar data by using a frequency-dependent model of the pulse profile.  This algorithm is a straightforward, yet novel extension of FFTFIT from~\\citet{Taylor92}, but it has some advantages over more standard techniques.  These include: \\begin{itemize} \\item{Simultaneous measurement of a phase (TOA) and dispersion measure (DM) across a channelized bandwidth, allowing for easy DM tracking.} \\item{{\\it In situ} accommodation for profile evolution, both intrinsic and extrinsic (eg. scattering), by the use of an arbitrary phase-frequency model (a ``portrait'') of the pulsar's signal.} \\item{Mitigation of scintillation effects by automatic weighting of frequency channels, which appropriately uses the information contained in the data.} \\item{Simplification of the timing procedure by eliminating (where appropriate) the use of phase offsets (i.e. JUMPs), multi-channel TOAs, and ad hoc methods for describing profile evolution, etc.} \\end{itemize} Any arbitrary model can be used, but the choice of model will affect the measured values.  We have made our code publicly available online.  The demonstration of a simple Gaussian modeling scheme to make these measurements in a 3-year, wideband dataset of the millisecond pulsar M28A shows that we are able to obtain reliable measurements of the dispersion measure, as well as improved TOA and DM precisions by up to a factor of four for the former and two for the latter.  The biggest improvements in the parameter precisions and in mitigating profile evolution were seen in the high signal-to-noise 1500~MHz data, which was our best ``wideband data'' for demonstration purposes in the sense that it has the largest fractional bandwidth (and therefore the most obvious effects from interstellar dispersion, scattering, and profile evolution).  We note that M28A was chosen for our demonstration precisely because it would show obvious improvements from using our new technique, and so its results may be atypical.  Similar improvements in other pulsars will depend on the mitigation of intrinsic and extrinsic profile evolution.  It became clear in our comparisons that there is a necessity for quantitative model selection based on more robust two-dimensional portrait modeling, which potentially can lead to the detection of a frequency-dependent DM, or other interesting signals.  We probed three typical SNR regimes with Monte Carlo tests and found that the algorithm performs well.  Except for the lowest SNR cases, the results from our Monte Carlo analyses has led us to the conclusion that a large number of frequency channels is appropriate for applying this technique.  A larger number of channels will provide the highest precision DM measurements and avoids averaging over profile evolution.  The proper incorporation of discrete DM measurements with their own heteroscedastic errors (besides the TOAs') into the determination of a timing model (eg. by using \\texttt{tempo}) is not trivial, but a Bayesian approach has been investigated in~\\citet{Lentati13}.  Relatedly, measuring DMs from non-simultaneous but temporally proximate multi-frequency data can be an intermediate improvement until larger bandwidths become readily available.  This is another avenue of future development, although it comes with the drawback of having correlated TOAs.  One important caveat in these Monte Carlo tests is that the model fitted to the simulated data was the true model from which the data were generated.  In practice, a Gaussian-component model fitted to real data will not match perfectly (leaving behind non-Gaussian residuals in all but the simplest or low SNR cases) and there will be a much stronger dependence of the measured DM and TOA on the number of channels and the amount of profile evolution.  This is another area requiring further testing, but it also suggests to err on the side of more frequency channels.  Although general, the algorithm will be most useful when applied to MSPs because of their sensitivities to small dispersion measure changes, as highlighted by Equation~\\ref{DMsmear}, and because of the need to correct for their profile evolution in wideband data to obtain the highest possible timing precisions.  For these reasons, we believe this algorithm will provide a natural TOA and DM measurement procedure for campaigns of MSP monitoring, like that of NANOGrav or other PTA experiments.  To determine how much is gained in timing precision (and, perhaps, sensitivity to gravitational waves), a direct application of our timing method to the wideband data from a subset of NANOGrav MSPs will be presented in a future paper."
    },
    "1402/1402.3731_arXiv.txt": {
        "abstract": "LUX, the world's largest dual-phase xenon time-projection chamber, with a fiducial target mass of 118~kg and 10,091 kg-days of exposure thus far, is currently the most sensitive direct dark matter search experiment. The initial null-result limit on the spin-independent WIMP-nucleon scattering cross-section was released in October~2013, with a primary scintillation threshold of 2~phe, roughly 3~\\keVnr~for LUX. The detector has been deployed at the Sanford Underground Research Facility (SURF) in Lead, South Dakota, and is the first experiment to achieve a limit on the WIMP cross-section lower than $10^{-45}$~cm$^{2}$. Here we present a more in-depth discussion of the novel energy scale employed to better understand the nuclear recoil light and charge yields, and of the calibration sources, including the new internal tritium source. We found the LUX data to be in conflict with low-mass WIMP signal interpretations of other results. ",
        "introduction": "The body of indirect evidence for dark matter is extensive. It includes the best-fit model for explaining the angular power spectrum of the Cosmic Microwave Background temperature anisotropy, gravitational lensing studies, large-scale structure observations and simulations, and galactic rotation curves~\\cite{Blumenthal:1984bp,Davis:1985,Clowe:2006,Hinshaw:2012aka,Ade:2013zuv}. All these point to a significant non-baryonic, cold (heavy and non-relativistic) component of matter in the Universe: $\\sim$85\\% of the matter, or, $\\sim$27\\% of the total mass-energy content of the Universe. The WIMP (Weakly Interacting Massive Particle) is a favored candidate, with many theories (for example, supersymmetry and Kaluza-Klein) providing natural candidate particles. Most direct dark matter searches are therefore geared towards finding WIMPs, which are expected to produce low-energy nuclear recoils (NR) in a detector, while electron recoils (ER) constitute a primary background to be reduced and identified~\\cite{PhysRevD.31.3059, Feng:2010}.  The noble element xenon has many favorable properties as a target for direct WIMP detection experiments~\\cite{araujoReview}. For such an element, deposited energy in the material is expressed in three possible channels: excitation, ionization, and heat. The most prominent channel for NR is heat, reducing the amount of energy in the first two significantly, in contrast to ER. Ionization electrons can recombine, or escape. Excitation and recombination together lead to the so-called primary scintillation signal (S1), with S1 from each source indistinguishable. Escaping ionization electrons lead to the secondary scintillation (S2) in the gas region of a two-phase detector. Lastly, noble elements are transparent to their own scintillation light because it originates in decaying molecules (excited dimers, i.e. excimers) rather than in atoms. ",
        "conclusions": "LUX has the largest exposure of any xenon TPC, as well as the lowest threshold. A new internal ER calibration source, tritium, was successfully implemented, and low-energy NR data agree with Monte Carlo, with the location of the NR band well predicted in terms of absolute charge and light yields for an electric field not studied previously. LUX has achieved the most stringent WIMP limit across a wide range of masses. In spite of assumptions more conservative than have been used in the past for xenon detectors, but which we have shown to agree well with the LUX data, our result is in conflict with low-mass WIMP interpretations of signals seen in DAMA, CoGeNT, CRESST, and CDMS Si. With a 300-day run LUX will probe cross-section versus mass parameter space previously unexplored by any other direct detection experiment, with a significantly improved sensitivity compared to this initial result."
    },
    "1402/1402.0075_arXiv.txt": {
        "abstract": "We present a new version of a semi-analytic model of cosmological galaxy formation, incorporating a star formation law with a feedback depending on the galaxy-scale mean dust opacity and metallicity, motivated by recent observations of star formation in nearby galaxies and theoretical considerations. This new model is used to investigate the effect of such a feedback on shaping the galaxy luminosity function and its evolution.  Star formation activity is significantly suppressed in dwarf galaxies by the new feedback effect, and the faint-end slope of local luminosity functions can be reproduced with a reasonable strength of supernova feedback, which is in contrast to the previous models that require a rather extreme strength of supernova feedback. Our model can also reproduce the early appearance of massive galaxies manifested in the bright-end of high redshift $K$-band luminosity functions.  Though some of the previous models also succeeded in reproducing this, they assumed a star formation law depending on the galaxy-scale dynamical time, which is not supported by observations.  We argue that the feedback depending on dust opacity (or metal column density) is essential, rather than that simply depending on gas column density, to get these results. ",
        "introduction": "\\label{sec:intro}  The basic picture of galaxy formation and evolution in the cosmological context can be explained in the standard $\\Lambda$ cold dark matter (CDM) cosmology. Particularly, large scale clustering properties and formation and evolution of dark matter halos can reliably be predicted by the theory of gravity.  However, in order to obtain the full picture of cosmological galaxy formation, we must solve complicated processes of baryonic physics, such as gas cooling, star formation, feedback, galaxy mergers, and so on.  One of the key observables about galaxies that must be explained by the theory of cosmological galaxy formation is the luminosity functions (LFs) and their evolution. Compared with the shape of dark matter halo mass function predicted by the $\\Lambda$CDM cosmology, the observed galaxy LFs have two remarkable features: flatter faint-end slopes and sharp exponential cut-off at the luminous/massive end (see Benson et al. 2003 and references therein), which must be explained by some baryonic processes.  A widely accepted solution to achieve a flat faint end is supernova feedback, i.e., energy input into the interstellar medium by supernova explosions to suppress star formation in small galaxies. However, the problem is not yet completely solved at the quantitative level. In fact, unreasonably high efficiency of supernova feedback to remove cold interstellar gas in dwarf galaxies is necessary in many existing theoretical models to reproduce the observed flat faint ends, and such an extreme supernova feedback tends to produce discrepancies with observations other than luminosity function shapes (Nagashima \\& Yoshii 2004, hereafter NY04; Nagashima et al. 2005; Bower et al. 2006, 2012; Guo et al. 2011; Wang, Weinmann \\& Neistein 2012; Mutch, Poole \\& Croton 2013; Puchwein \\& Springel 2013; Hopkins et al. 2013). These results imply that another physical effect may also be taking an important role to produce the observed flat faint end slopes.  For the massive end, a popular solution to suppress the formation of too massive galaxies is the feedback by active galactic nuclei (AGNs; e.g., Bower et al. 2006; Croton et al. 2006; Menci et al. 2008; Somerville et al. 2008; Guo et al. 2011).  The AGN feedback can also explain the observed trends of the early appearance of massive and quiescent galaxies at high redshifts, and downsizing of star-forming galaxies from high to low redshifts, which are apparently in contradiction with the simple expectation in the $\\Lambda$CDM universe (e.g., Bower et al. 2006; Somerville et al. 2008). However, there are large uncertainties about the physics of AGN feedback both in theoretically and observationally. The current success in explaining the observed trends by this process is based on rather phenomenological modelings including highly uncertain parameters, and further studies are required to confirm the quantitative influence of this process on galaxy evolution.  Therefore it is still worth to explore yet other physical effects working to shape galaxy LFs, which is the aim of this paper. It is reasonable to expect that such an effect would be manifested in the scaling laws about star formation efficiency.  The relation between the surface densities of star formation rate (SFR) and gas ($\\Sigma_{\\rm SFR}$-$\\Sigma_{\\rm gas}$) has been a subject of intensive research. It is popular to fit this relation by a power-law (so-called Kennicutt-Schmidt law, Kennicutt 1998), but recent observations indicate a cut-off around the total (i.e., {\\hbox{H{\\sc i} }} + H$_2$) gas density of $\\Sigma_{\\rm gas} \\sim 10\\;M_\\odot \\ \\rm pc^{-2}$, under which SFR is suppressed and not well correlated with gas density.  This threshold gas density for SFR can be interpreted as a result of less efficient formation of cold molecular gas under the threshold, while the star formation efficiency (SFE) from molecular gas is rather universal in many different environments (Wong \\& Blitz 2002; Kennicutt et al. 2007; Bigiel et al. 2008, 2010; Leroy et al. 2008; Blanc et al. 2009; Heiderman et al. 2010; Lada et al. 2010, Schruba et al. 2011; see Kennicutt \\& Evans 2012 and Schruba 2013 for reviews).  A likely physical origin of the suppression of H$_2$ formation under the threshold is radiative feedback by UV photons produced by young massive stars (Schaye 2004; Krumholz, Mckee \\& Tumlinson 2008, 2009; McKee \\& Krumholz 2010; Hopkins et al. 2013).  The formation of H$_2$ is driven by collisionally excited metal line cooling and molecule formation on dust grain surfaces, which should be balanced with molecule dissociations by UV photons and grain photoelectric heating, both of which are energetically supplied by UV radiation field. If a region in a galaxy is optically thick to UV radiation field by dust grains, self-shielding of UV radiation would accelerate H$_2$ formation. This implies that the more fundamental threshold about star formation is not the total gas surface density but dust opacity.  For a typical dust-to-gas ratio, the observationally indicated threshold in $\\Sigma_{\\rm gas}$ is close to the value at which the effective dust opacity $\\tau_d^{\\rm eff}$ becomes of order unity, where $\\tau_d^{\\rm eff}$ is averaged over wavelength with a weight of the heating radiation energy spectrum (Totani et al. 2011).  Therefore it is physically reasonable to expect that a galaxy-scale mean value of $\\tau_d^{\\rm eff}$ has an important role in galaxy formation and evolution.  A further observational support to this picture comes from infrared observations. The relations between dust temperature, galaxy size, and infrared luminosity of $\\sim$1,000 nearby star-forming galaxies indicate that almost all of them are in the optically thick regime, and the distribution of dust opacity estimated by gas-phase metal column density suddenly drops around $\\tau_d^{\\rm eff} \\sim 1$, indicating less efficient formation of galaxies at $\\tau_d^{\\rm eff} \\la 1$ (Totani et al. 2011).  The purpose of this paper is to investigate the effect of the radiative feedback depend on dust opacity, on cosmological galaxy formation and evolution particularly about the shape of galaxy luminosity functions.  The theory of structure formation in the universe predicts that the mean surface density $M/r^2$ of dark halos with mass $M$ and size $r$ nearly scales as $\\propto M^{1/3} (1+z)^2$, indicating higher gas surface density and dust opacity at higher redshifts in more massive objects, and hence more efficient star formation.  This may have a favorable effect to explain observations, in a similar way to the feedbacks by supernovae and AGNs.  To investigate the effect quantitatively, we use a semi-analytic model (SAM) of cosmological galaxy formation, {\\it the Mitaka model} (NY04). This is a model similar to general SAMs, in which formation and evolution of dark matter halos are solved analytically or calculated by N-body simulations, while complicated baryonic processes are treated phenomenologically (for reviews, see Baugh 2006; Benson 2010). In general, SAMs has many adjustable parameters and the effects of complicated physical processes on the LFs are degenerate (e.g., Neistein \\& Weinmann 2010); therefore a set of best-fit parameters may not be a quantitatively correct description of real galaxy formation. It should be noted that the most important aim of this work is to examine the qualitative effects of the new feedback on luminosity functions.  In most of the SAMs, the star formation rate is simply proportional to cold gas mass, and the star formation time scale is modeled as a simple function of the dynamical time scale of galaxy disks or DM halos (e.g., Cole et al. 2000; NY04).  Some models (e.g., Kauffmann 1996; Croton et al. 2006; Somerville et al. 2008; Lagos et al. 2011; Wang et al. 2012) incorporated the threshold of gas surface density below which star formation activity is significantly suppressed.  In the models of Kauffmann (1996), Croton et al. (2006), and Lagos et al. (2011), they introduced the threshold of gas surface density motivated by the Toomre stability criterion on a galactic scale (Toomre 1964).  In this scenario the threshold of gas surface density increases with redshift, and hence the threshold effect should be systematically different in the cosmological context from the threshold by dust opacity considered in this paper. Furthermore, some recent observations indicate that star formation are controlled by the physical state of local interstellar gas, rather than the dynamical state of an entire galaxy (e.g., Leroy et al. 2008; Lada et al. 2010).  In other models, such as Somerville et al. (2008), a critical gas surface density threshold for star formation is introduced motivated from the observations of the $\\Sigma_{\\rm SFR}$-$\\Sigma_{\\rm gas}$ relation; however, to our knowledge there are no SAMs that consider a feedback depending on dust surface density rather than gas density. Recently Krumholz \\& Dekel (2012) incorporated a star formation law which depends on gas surface density and gas metallicity, and discussed average evolution of typical galaxies without calculating detailed merger histories of dark halos. The relation between the luminosity function shapes and the dust opacity threshold of star formation has not yet been discussed in previous studies.  This paper is organized as follows.  In section \\ref{sec:model}, we will describe our model particularly focusing on the modelings of star formation and feedback.  In section \\ref{sec:Results}, we show the results of our model, and section \\ref{sec:gasmodel} is devoted for discussion.  We will summarize our work in section \\ref{sec:summary}. In this work, the cosmological parameters of $\\Omega_{0}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_{0}=70~{\\rm Mpc^{-1}~km~s^{-1}}$ are adopted, and all magnitudes are expressed in the AB system.   \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=0.49\\hsize]{g_LF_NY04_z0p0_paper.ps} & \\includegraphics[width=0.49\\hsize]{K_LF_NY04_z0p0_paper.ps} \\\\ \\end{tabular} \\caption{ ({\\it Left}) The local $g$-band LF compared with the NY04 model. The solid and dashed lines represent the NY04 model with strong SN feedback ($\\alpha_{\\rm hot} = 4$) and weak (reasonable) SN feedback ($\\alpha_{\\rm hot} = 2$), respectively. We only show the results of CSF model, since the DSF model gives the almost same results. Filled circles indicate the SDSS $g$-band LF obtained by Blanton et al. (2005), and open circles are the 6dF $b_{\\rm j}$-band LF obtained by Jones et al. (2006). We have transformed the 6dF $b_{\\rm j}$-band LF to match $g$, by subtracting 0.25 mag (Blanton et al. 2005). ({\\it Right}) The same as the left panel but for the local $K$-band LF.  Data points are the 6dF galaxy survey (Jones et al. 2006) and 2MASS (Kochanek et al. 2001). } \\label{fig:NY04_LF_z0p0} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  In this paper, we have considered a new feedback mechanism on star formation depending on galaxy-scale mean optical depth to absorption by dust grains, and examined the effect on galaxy luminosity functions and their cosmological evolution, making use of a semi-analytic model of galaxy formation.  The introduction of such feedback process is motivated not only by theoretical considerations but also by recent observations, which indicate that star formation activity is significantly suppressed in galaxies that are transparent to UV radiation.  The structure formation theory predicts that the dust-opacity becomes higher in massive objects and at higher redshifts for a fixed dust-to-gas ratio; therefore it is expected that the faint-end of local LFs would be suppressed, which is required for the current galaxy formation models to match the observations. Note that extremely strong supernova feedback was required in the conventional models to reproduce the observed faint end of local LFs. Such feedback process would also accelerate the formation of massive galaxies at high redshifts.  We have tested a few models about star formation feedback, and the best fit with observations is found with the model in which star formation is suppressed when the galaxy-scale dust opacity is low and metallicity is higher than a critical value (the evolving $\\varepsilon_{\\min}$ model).  The latter condition is introduced phenomenologically, but theoretically motivated by the process of photoelectric heating by dust grains. In this model formation of dwarf galaxies at $z \\sim 0$ is significantly suppressed, and the model successfully reproduces the faint-end slope of local LFs with a physically natural strength of the SN feedback.  The new model also succeeded in reproducing the number density of high-$z$ massive galaxies.  The early appearance of massive galaxies have been explained by the AGN feedback process; however we have found that the star formation model is also important as well as the AGN feedback.  In most of SAMs, star formation time scale is assumed to be proportional to the dynamical time scale of a host DM halo or galaxy disk (e.g., Cole et al. 2000; Bower et al. 2006; NY04). This is essential to explain the early appearance of massive galaxies, because the model with a constant star formation time scale cannot reproduce it even if the AGN feedback is incorporated.  However, recent observations suggest that the star formation efficiency is closely related to the gas or dust surface density, rather than the dynamical time scale of an entire galaxy or halo (see section \\ref{sec:intro}). Our new model incorporating the AGN feedback can explain the number density of high-$z$ massive galaxies with the observationally suggested star formation law. The new model is also consistent with the observed cosmic star formation history.  We also tested a star formation feedback model depending simply on the gas surface density (the $\\Sigma_{\\rm gas}$ model), rather than the dust opacity, to examine whether the dust opacity is essential or not. Although this model can also reproduce the shape of the local LFs, the difference from the evolving $\\varepsilon_{\\min}$ model appears in the mass function (or $K$-band LF) at high redshifts.  The evolving $\\varepsilon_{\\min}$ model predicts more galaxies than the $\\Sigma_{\\rm gas}$ model at the bright end of $K$-band LFs at $z \\sim 2$, which is in better agreement with the observed data.  To conclude, we have found that the feedback depending on galaxy-scale dust opacity has significant effects on the cosmological galaxy formation, and has good properties to solve some of the problems found in the previous theoretical models.  However, it should also be noted that there are still various uncertainties in our model.  For example, we determined the value of star formation efficiency under the dust opacity threshold phenomenologically from fits to the luminosity function data, but these results should be examined in light of theoretical studies of star formation.  We assumed that dust mass is simply proportional to the metal mass, but it is not obvious that this proportionality is valid for all galaxies. More observational and theoretical studies on formation/evolution of dust grains are desirable to establish a better star formation modeling for cosmological galaxy formation."
    },
    "1402/1402.0596_arXiv.txt": {
        "abstract": "{ Recently, there are several reports that the cosmic magnetic fields on Mpc scale in void region is larger than $\\sim 10^{-15}$G with an uncertainty of a few orders from the current blazar observations. On the other hand, in inflationary magnetogenesis models, additional primordial curvature perturbations are inevitably produced from iso-curvature perturbations due to generated electromagnetic fields. We explore such induced curvature perturbations in a model independent way and obtained a severe upper bound for the energy scale of inflation from the observed cosmic magnetic fields and the observed amplitude of the curvature perturbation , as $\\rho_\\inf^{1/4}< 30{\\rm GeV}\\times(B_\\obs/10^{-15}{\\rm G})^{-1}$ where $B_\\obs$ is the strength of the magnetic field at present. Therefore, without a dedicated low energy inflation model or an additional amplification of magnetic fields after inflation, inflationary magnetogenesis on Mpc scale is generally incompatible with CMB observations. } ",
        "introduction": "It has been known for a long time that galaxies and galactic clusters have their own magnetic fields~\\cite{Wielebinski:2005,Bernet:2008qp,Bonafede:2010xg,Beck:2012}. However, their origin is a big mystery of astronomy and cosmology~\\cite{Giovannini:2003yn,Kandus:2010nw,Durrer:2013pga}. Recently the generation mechanism of the magnetic fields in the universe attracts much attention because there are several reports that magnetic fields are found even in void regions. Such void magnetic fields could be detected by blazar observations~\\cite{Neronov:1900zz, Tavecchio:2010mk, Essey:2010nd, Taylor:2011bn, Takahashi:2013uoa, Finke:2013bua} and it is reported that their strength is larger than $\\sim 10^{-15}$G with an uncertainty of a few orders. On the other hand, the upper bound on primordial magnetic fields could be also obtained from the cosmic microwave background (CMB) and the large scale structure (LSS) observations, and current upper bound is roughly given by $10^{-9}$G (see, e.g. \\cite{Yamazaki:2012pg,Shaw:2010ea}, and references therein) \\footnote{ Ref. \\cite{Kawasaki:2012va} reported an updated constraint on a primordial magnetic field during big bang nucleosynthesis (BBN) as $10^{-6} $G.}. Therefore we know there exist the magnetic fields in the universe with the strength, \\footnote{ The upper bound is irrelevant for magnetic fields which are produced after CMB photons are radiated.} \\begin{equation} 10^{-15}\\G \\lesssim B_\\obs \\lesssim 10^{-9}\\G. \\label{Allowed range} \\end{equation} Nevertheless, their origin is still unknown and no successful quantitative model is established. If the magnetic fields are produced in the primordial universe, they can seed the observed galactic and cluster magnetic field~\\cite{Ryu:2008hi,Pakmor:2013rqa} as well as directly explain the void magnetic fields.   As one of the mechanism of generating such cosmic magnetic fields, ``inflationary magnetogenesis\" has been widely discussed. In the context of the inflationary magnetogenesis, large scale magnetic fields, as well as the primordial curvature perturbations, are basically generated from the quantum fluctuations. Although many models of the generation of magnetic fields during inflation are proposed so far~\\cite{Turner:1987bw, Ratra:1991bn, Gasperini:1995dh, Davis:2000zp, Bassett:2000aw, Bamba:2003av,Enqvist:2004yy, Martin:2007ue, Ferreira:2013sqa,Garretson:1992vt,Anber:2006xt}, it is known that these inflationary magnetogenesis models suffer from several problems, namely the strong coupling problem~\\cite{Gasperini:1995dh,Demozzi:2009fu,Fujita:2012rb}, the backreaction problem~\\cite{Bamba:2003av,Demozzi:2009fu,Kanno:2009ei,Fujita:2012rb}, the anisotropy problem~\\cite{Bartolo:2012sd,Thorsrud:2013kya} and the curvature perturbation problem~\\cite{Suyama:2012wh,Barnaby:2012xt,Giovannini:2013rme,Ringeval:2013hfa,Fujita:2013qxa,Nurmi:2013gpa}. In particular, the curvature perturbation problem, where the primordial curvature perturbations which are induced from the generated electromagnetic fields during inflation should not exceed the observed value of CMB experiments, gives strong constraints on inflationary magnetogenesis models. For examples, in our previous paper~\\cite{Fujita:2013qxa}, we have intensively studied the curvature perturbation problem by using a specific model, so-called the kinetic coupling model~\\cite{Ratra:1991bn}, and showed that the allowed strength of the produced magnetic fields is far weaker than the observational lower bound given by eq.~\\eqref{Allowed range}. Ref.~\\cite{Ringeval:2013hfa} have investigated the curvature perturbation problem specifying the time evolution of the magnetic fields during inflation as the power-law of the conformal time and showed limits of the amplitude of the present magnetic fields for the monomial and the hill-top inflation models with several reheating scenarios.  Although investigation of the constraint on inflationary magnetogenesis in model dependent ways is important, to discuss whether inflationary magnetogenesis is {\\it really} possible or not, model independent arguments should be also necessary. As for such discussion, in ref.~\\cite{Suyama:2012wh} the authors have put the lower bound on the inflation energy scale $\\rho_{\\rm inf}$ only by requiring the production of magnetic fields with the sufficient strength $B_{\\rm obs} \\sim 10^{-15}$G, but they assumed that the dominant primordial curvature perturbation is generated during the single slow-roll inflation. In ref.~\\cite{Fujita:2012rb}, apart from the curvature perturbation problem, by requiring to escape from the strong coupling and the backreaction problems, the upper bound on $\\rho_{\\rm inf}$ has been % put in model independent ways.  In this paper, we consider the curvature perturbation problem of inflationary magnetogenesis % in a model independent way and we do not specify the dominant contribution of the primordial curvature perturbations. That is, our result could be also applied to the case where the dominant primordial curvature perturbation is sourced from a light scalar field other than inflaton. We focus on the existence of the electric fields due to the time evolution of the magnetic fields in the Friedmann-Lemaitre-Robertson-Walker (FLRW) universe and we show that if one requires inflation magnetogenesis is responsible for the generation of the observed magnetic fields and assumes no additional amplification after inflation, the inflation energy scale is constrained by the curvature power spectrum $\\mathcal{P}_\\zeta$ as \\begin{equation} \\mathcal{P}_\\zeta^\\obs > \\mathcal{P}_\\zeta^\\em \\quad\\Rightarrow\\quad \\rho_{\\rm inf}^{1/4} < 30{\\rm GeV}\\times \\left(\\frac{p_B}{1{\\rm Mpc}^{-1}}\\right)^{\\frac{5}{4}} \\left( \\frac{B_{\\rm obs}}{10^{-15}{\\rm G}} \\right)^{-1}, \\label{result intro} \\end{equation} where $\\rho_\\inf$ is the energy scale of inflation, $p_B > 1{\\rm Mpc}^{-1}$ is the peak wave number of the void magnetic field and $B_\\obs$ is the magnetic field strength today. Therefore, our result indicates some tension between inflationary magnetogenesis and phenomenologies in the very early universe, e.g., genesis of the baryon or dark matter, where high energy physics are involved.  We also discuss a possible way out of our constraint. If strong magnetic fields are produced without amplifying electric fields, one could avoid our constraint. Such situation is apparently realized in a tree-level analysis of the so-called strong coupling regime of the kinetic coupling model ~\\cite{Ratra:1991bn, Gasperini:1995dh, Bamba:2003av}. However, since the coupling constant becomes huge in the model, a non-perturbative analysis beyond the tree-level is required to make the correct prediction~\\cite{Demozzi:2009fu}. Furthermore, an additional amplification or a non-adiabatic dilution of magnetic fields after inflation can relax our constraint. For example, if the inverse cascade works, the constraint is alleviated~\\cite{Campanelli:2007tc, Saveliev:2013uva}.   The rest of paper is organized as follows. In section 2, we briefly review the current lower bound on the cosmic magnetic field from the blazar observations and outline how we constrain inflationary magnetogenesis in a model independent way. In section 3, we derive an expression of curvature perturbations induced by electromagnetic fields during inflation. In section 4, the constraint on inflationary magnetogenesis is obtained. Section 5 is devoted to a summary and discussions. In appendix, we discuss the constraint without the assumption of the instantaneous reheating. ",
        "conclusions": "In this paper, we show that inflationary magnetogenesis is generally constrained as eq.~\\eqref{final expression} by requiring that the curvature perturbation induced by the electric field during inflation should be smaller than the Planck observation value:  $\\mcP_\\zeta^\\obs(k) =2.2\\times10^{-9}$. We emphasize that our argument is model independent as we outlined in sec.~\\ref{Model independent approach}. The main result eq.~\\eqref{final expression} indicates that inflationary magnetogenesis is under pressure in several ways.  First, it is known that the reheating (thermalization) energy scale is bounded as $\\rho_{\\rm reh}^{1/4} \\gtrsim 10$MeV in order to achieve a successful BBN~\\cite{Hannestad:2004px}. Therefore even if eq.~\\eqref{final expression} is almost saturated, for example $\\rho_{\\inf}^{1/4}\\sim 10$GeV, the reheating should be quickly completed.  Second, the generation of the observed curvature perturbation is in danger. Eq.~\\eqref{final expression} can be translated as \\begin{equation} H_\\inf< 2\\times 10^{-7} {\\rm eV} \\left(\\frac{p_B}{1{\\rm Mpc}^{-1}} \\right)^{5/2} \\left(\\frac{B_\\obs}{10^{-15}\\G} \\right)^{-2}, \\label{H constraint} \\end{equation} where $H_\\inf$ is the Hubble parameter during inflation. For a scalar field to acquire a perturbation during inflation, its mass should be smaller than $H_\\inf$. Thus inflaton field or a spectator field which is responsible to produce $\\mcP_\\zeta^\\obs$ must be extremely light during inflation. During reheating era, however, it has to quickly decay into the standard model particles to cause the BBN properly. Furthermore, in the case of single slow roll inflation, eq.\\eqref{H constraint} and the COBE normalization indicate an extreme slow-roll $\\epsilon< 4\\times 10^{-62}$ which demands a dedicated inflation model. It is interesting to note that eq.\\eqref{H constraint} corresponds to the very small tensor-to-scalar ratio, $r< 7\\times 10^{-61}$. Hence a detection of background gravitational waves in the future excludes inflationary magnetogenesis. \\footnote{See ref.~\\cite{Ferreira2} in which our model-independent constraint is followed up in the light of the BICEP2 result~\\cite{Ade:2014xna}. }   Third, in such a low reheating temperature, thermal production of the dark matter or the baryon number seems hopeless. Since $30$GeV is accessible by colliders, effects beyond the standard model have been severely restricted. To realize the dark matter production and baryogenesis, a non-thermal mechanism like the direct decay of inflaton should be considered.  In spite of these negative implications, since we have the observational evidence of the magnetic fields in the universe and we are lack of a plausible magnetogenesis model, the inflationary origin of the magnetic field is still an appealing idea. It should be noted that we assume no amplification of the magnetic fields after inflation to derive eq.~\\eqref{final expression}. Thus our result might imply that inflationary magnetogenesis need an additional amplification or a non-adiabatic dilution of magnetic fields after inflation. If the magnetic field generated during inflation is amplified by some mechanism like preheating process~\\cite{Finelli:2000sh} or the inverse cascade~\\cite{Campanelli:2007tc, Saveliev:2013uva}, the constraint is alleviated.  Another possible way out from our constraint is to produce a large amplitude of the vector potential before the horizon crossing. It is known that, in the so-called strong coupling regime of the kinetic coupling model, the electric field is not much stronger than the magnetic field and the backreaction and curvature perturbation problems are evaded (if loop effects are neglected)~\\cite{Demozzi:2009fu}. This is because the vector potential $\\mcA_k$ is almost constant on super-horizon ($n\\simeq 0$ in our language). The magnetic field is produced since $\\mcA_k$ has a large amplitude at the horizon crossing due to the small kinetic function. However, as discussed below eq.~\\eqref{def of rho^em}, such model suffers from the strong coupling problem and reliable calculations are difficult to be done. If a large amplitude of a static vector potential is realized without the strong coupling or one can take into account the loop effects in some non-perturbative way, sufficient magnetogenesis might be achieved."
    },
    "1402/1402.7180_arXiv.txt": {
        "abstract": "The advent of synoptic sky surveys has spurred the development of techniques for real-time classification of astronomical sources in order to ensure timely follow-up with appropriate instruments.  Previous work has focused on algorithm selection or improved light curve representations, and naively convert light curves into structured feature sets without regard for the time span or phase of the light curves.  In this paper, we highlight the violation of a fundamental machine learning assumption that occurs when archival light curves with long observational time spans are used to train classifiers that are applied to light curves with fewer observations. We propose two solutions to deal with the mismatch in the time spans of training and test light curves.  The first is the use of classifier committees where each classifier is trained on light curves of different observational time spans.  Only the committee member whose training set matches the test light curve time span is invoked for classification.  The second solution uses hierarchical classifiers that are able to predict source types both individually and by sub-group, so that the user can trade-off an earlier, more robust classification with classification granularity.   We test both methods using light curves from the MACHO survey, and demonstrate their usefulness in improving performance over similar methods that naively train on all available archival data. ",
        "introduction": "The advent of next generation optical and radio telescopes such as the Large Synoptic Survey Telescope (LSST) \\cite{tyson_large_2003} and the Square Kilometre Array (SKA) \\cite{taylor2013} will enable massive wide-field surveys for highly variable or transient astronomical sources. Prompt multi-wavelength follow-up of candidate sources will be critical in achieving the scientific goals of these surveys. A global network of potential follow-up instruments is available, but due to limited resources, only high quality candidates will be selected for follow-up. In the regime of big data surveys, manual inspection of data is no longer possible, therefore automated classification methods must form part of the data pipelines to determine whether a detected candidate is an object of interest that requires follow-up.  The use of machine learning for source type classification of archival light curve data is well-established \\cite{eyer_automated_2004,debosscher_automated_2007,sarro_automated_2008-1,kim2011,richards2011,blomme_automated_2011}. Some methods focus on specific events or source types to classify \\cite{kim2011}, while others focus on periodic variables \\cite{blomme_automated_2011} or a range of source types \\cite{debosscher_automated_2007,richards2011}.  All must contend with the challenges inherent in classifying light curve data including the generation of robust training data \\cite{long_optimizing_2012}, the search for discriminating features \\cite{debosscher_automated_2007,richards2011} and methods that are robust to unevenly sampled data, as well as missing, spurious and sparse observations. Various classification algorithms, such as neural networks, support vector machines and random forests, have been explored. However, algorithmic selection is not the most significant factor in achieving high performance.  Current surveys that use machine learned classification in their data processing pipelines include the Palomar Transient Factory (PTF) \\cite{law_palomar_2009} and the Catalina Real-time Survey (CRTS) \\cite{djorgovski_catalina_2011}. PTF employs an automated ``real-bogus'' classification system that identifies true astronomical transient candidates using features extracted from candidates in subtracted images \\cite{bloom_towards_2008, brink_using_2012}. To date, the PTF system performs binary classification (real or not), and primarily uses image features rather than light curve information in its decision-making. CRTS has developed a real-time transient detection pipeline that contends with sparse and heterogeneous data sources \\cite{djorgovski_towards_2011}, self-updates as observations are received, and makes robust decisions on known classes while potentially discovering unknown sources.  Whether performing real-time or archival classification, the use of machine learning algorithms in the methods cited above necessitates the creation of a structured data set of features from light curves that are arbitrary in time span and out of phase with each other.  After training a classifier on this set, a partially-observed light curve is converted into the same feature characterization and classified.  Theoretically the test examples belong to the same population as the training examples, however the fact that the test light curves are only partially-observed implies that the extracted features may effectively belong to a different statistical distribution.  This violates a basic premise of all supervised learning algorithms that both training and test examples must belong to the same distribution.  A large enough violation of this premise will degrade the performance of the classifier, as it cannot generalize to new test data. To our knowledge, this issue has not  been addressed in previous work.   This paper discusses how to do classification robustly with light curve data arriving in a stream. This type of algorithm is known as an {\\it online learning algorithm} and it incorporates new data in the classification model as they become available. There are two novel aspects to our work.  First, in order to handle the mismatch between the light curve time spans in training and test sets, we developed a method that builds a committee of classifiers each trained with light curves of a set time span.  The appropriate committee member that matches the test light in time span is invoked for classification.  We show that this method that can outperform a naive method that trains a single classifier using all available archival data.  Second, we implemented hierarchical classifiers. A non-hierarchical classifier assigns a single label from a set of discrete labels to a test example, while a hierarchical classifier assigns either a single label or a sub-grouping of discrete labels, where sub-groups are pre-defined by the end user in a tree-like structure.  The advantage of a hierarchical classifier is its ability to trade classification granularity against earlier predictions. Given a user-specified desired confidence level, the classifier may not be able to predict down to a single discrete label (e.g., RR Lyrae) but able to predict an informative sub-group of the classes (e.g., variable stars).  This enables the end user of the system to receive some information earlier rather than a more confident prediction much later. ",
        "conclusions": "\\label{s_discussion}  In this paper, we explored the feasibility of using light curves for classifying transient and variable sources in an online setting. Although machine-learned classification with light curves has been used in astronomy, classification with a {\\em data stream} (online classification) has not received much attention. As we approach the era of large synoptic surveys, online classification will become increasingly relevant. Transient surveys will benefit from having a classification module as part of the processing pipeline.   \\begin{table}[h!] \\centering \\begin{tabular}{ lrrrrr } \\hline &  Oracle &  \\multicolumn{2}{c}{Naive} &  \\multicolumn{2}{c}{Committee }\\\\ & ${\\sim}$2500 days& 160 days & 370 days & 160 days & 370 days \\\\ \\hline \\hline Non-variables & $99\\%$ &  $75\\%$ & $96\\%$ &  $98\\%$  & $99\\%$  \\\\ Variables        & $91\\%$  & $67\\%$ & $80\\%$ &   $81\\%$  & $86\\%$ \\\\ \\hline Overall           & $97\\%$ & $69\\%$ & $91\\%$  &   $93\\%$  & $95\\%$ \\\\ \\hline  \\end{tabular} \\caption{Comparison of classification accuracies using different methods} \\label{tab:comp} \\end{table}  \\begin{table}[h!] \\centering \\begin{tabular}{ lrrrr } \\hline & \\multicolumn{2}{c}{Hierarch} & \\multicolumn{2}{c}{Hierarch + threshold of 0.7}   \\\\ & 160 days & 370 days & 160 days & 370 days \\\\ \\hline \\hline Non-variables     & $97\\%$ & $98\\%$ & $98\\%$ (95\\% of srcs) &$99\\%$ (96\\% of srcs) \\\\ Variables\t\t   & $82\\%$ &  $88\\%$ & $99\\%$ (68\\% of srcs) &$99\\%$ (77\\% of srcs) \\\\ \\hline Overall               &  $93\\%$ &  $95\\%$ & $98\\%$ & $99\\%$ \\\\ \\hline  \\end{tabular} \\caption{Comparison of leaf node classification accuracies using hierarchical classification.} \\label{tab:hierarch} \\end{table}  An online classifier that consists of a committee of classifiers, each learned from light curves of different time spans that match that of the source to be classified, performs significantly better than a classifier that learns from a fully observed light curve. This is shown in the results summarized in Table \\ref{tab:comp}. However, the performance differs dramatically across source types. This shows that online classification with light curves has potential for some source classes, but will need to be supplemented with other information to be useful for other source classes.  We also explored using hierarchical classification with the MACHO dataset. By stopping the classification at the level of the hierarchy where the classification probability is less than the confidence threshold, we can achieve a higher level of accuracy but at a coarser level of classification. The results are summarised in Table \\ref{tab:hierarch}. One area to explore is to evaluate the performance of a hierarchical classifier where the hierarchy structure matches that of the follow-up observation requirements. The result of such a classifier would thus be meaningful and can be used for prioritising follow-up observations. More work can also be done on optimising the classifier at each node by selecting the best feature set and algorithm."
    },
    "1402/1402.2246_arXiv.txt": {
        "abstract": "{A middle-aged radio-quiet pulsar J0357+3205 was discovered in gamma-rays with \\textit{Fermi} and later in X-rays with \\textit{Chandra} and \\textit{XMM-Newton} observatories. It produces an unusual thermally-emitting pulsar wind nebula observed in X-rays. } { Deep optical observations were obtained to search for the pulsar optical counterpart and its nebula using the Gran Telescopio Canarias (GTC). } { The direct imaging mode  in the Sloan $g'$ band was used. Archival X-ray data were reanalysed and compared with the optical data. } { No pulsar optical counterpart was detected  down to $g'$~$\\geqslant$ 28\\fm1.  No  pulsar nebula was either identified in the optical.  We confirm early results that the X-ray spectrum of the pulsar consists of a nonthermal power-law component of the pulsar magnetospheric origin dominating at high energies and a soft thermal component from the neutron star surface.  Using  magnetised partially ionised hydrogen atmosphere models in X-ray spectral fits we found that the thermal component can come from entire surface of the cooling neutron star with a temperature of 36$^{+9}_{-6}$ eV, making it one of the coldest among cooling neutron stars known. The surface temperature agrees with the standard neutron star cooling scenario. The optical upper limit does not put any additional constraints on the thermal component, however it implies  a strong spectral break for the nonthermal component between the optical and X-rays as is observed in other middle-aged pulsars. } { The thermal emission from the entire surface of the neutron star likely dominates over the nonthermal emission in the UV range. Observations of the \\psr~in this range are promising to put more stringent constraints on its thermal properties. } ",
        "introduction": "\\label{sec1} Gamma-ray pulsars are considered as one of the main targets of the \\textit{Fermi} mission. For the five years of activity The Large Area Telescope (LAT) has discovered numerous amounts of such sources previously observed in the radio band. But, apart from the ability to detect many known radio pulsars in $\\gamma$-rays, \\textit{Fermi} LAT also affords the opportunity to discover pulsars independently in so-called blind searches \\citep[cf.][]{Saz09}. These blind searches were quite successful leading to discovery of about three dozens of pulsars in $\\gamma$-rays \\citep[see e.g.,][]{Saz13, Pletsch}. Further multiwavelength investigations of these objects are crucial for unveiling the pulsar emission nature. Because \\textit{Fermi} pulsars are typically nearby and energetic \\citep{Saz13}, they, in particular, appear to be promising targets for studies in X-ray and optical domains.  A middle-aged radio-quiet PSR J0357$+$3205 with period $P$~= 444~ms, magnetic field $B$~= 2.3~$\\times$~10$^{12}$~G and characteristic age $P$/2$\\dot{P}$~= 5.4~$\\times$ 10$^{5}$ yr was discovered in one of the \\textit{Fermi} LAT blind frequency searches \\citep{BlindFreqAbdo2009}. The distance to the pulsar of about 500 pc was estimated by \\citet{XrayFirstDeLuca} based on the $\\gamma$-ray ``pseudo-distance'' relation \\citep[see e.g.,][]{sazparkinson2010ApJ}. First X-ray observations of the pulsar field with \\textit{Chandra} had revealed a faint X-ray counterpart of the object with an extended (9 arcmin) X-ray tail \\citep{XrayFirstDeLuca}. Subsequent \\textit{XMM-Newton} observations had shown clearly that emission from the pulsar itself is generally nonthermal with a soft thermal component \\citep{XraySecondDeLuca}. The pulsar field was also observed in the optical and near-infrared bands with 2.5--4 m class telescopes. No counterpart was found down to $V \\ga $ 26\\fm7 \\citep{XrayFirstDeLuca}.  To search for an optical counterpart of PSR J0357$+$3205 and/or its tail at a higher sensitivity level, we performed deep optical observations  with the 10.4 m GTC. The details of  observations and data reduction are described in Sect.~\\ref{sec2}, our results together with reanalysis of the archival X-ray data are presented in Sect.~\\ref{sec3} and are discussed in Sect.~\\ref{sec4}.  \\begin{figure*}[t] \\setlength{\\unitlength}{1mm} \\begin{center} \\begin{picture}(145,65)(0,0) \\put (-20,-20) {\\includegraphics[width=100.0mm, clip=]{paper_2.ps}} \\put (80,-12.5) {\\includegraphics[width=80mm, clip=]{small.ps}} \\end{picture} \\end{center} \\vspace{21mm} \\caption{{\\sl Left} panel: GTC/OSIRIS $\\sim$~51\\asec~$\\times$~51\\asec~Sloan $g'$-image fragment of the \\psr~field. The circle shows 3$\\sigma$ X-ray pulsar position uncertainty for the optical observations epoch (see text for details). The 15\\asec~$\\times$~15\\asec~pulsar vicinity within the dashed rectangle is enlarged in the {\\sl right} panel and smoothed with a one pixel Gaussian kernel. The sources discussed in the text are labelled by numbers. } \\label{fig:1} \\end{figure*} ",
        "conclusions": "\\begin{figure}[t] \\begin{center} \\setlength{\\unitlength}{1mm} \\resizebox{15.5cm}{!}{ \\begin{picture}(50,24.5)(0,0) \\put (-1,0) {\\includegraphics[scale=0.127]{0357_superspec12.eps}} \\end{picture} } \\end{center} \\caption{ Unabsorbed spectrum of PSR J0357$+$3205 in the optical-X-ray range. The GTC $g'$-band flux upper limit is indicated by the bar with an arrow. The \\textit{XMM-Newton} and \\textit{Chandra} data points are shown by grey bars.\\protect\\footnotemark\\ The thick solid line is the best-fit BB+PL model to the X-ray data. The thin dotted line with the dark-grey filled region is the PL component with 90\\% uncertainties. The thick dotted line is the BB hot spot component. The thin/thick dash-dotted and dashed lines are the best-fit NSA and NSMAX components with variable/fixed normalisations. The upper limit on the thermal spectral flux from the entire surface of the NS obtained with the BB+PL+BB model is shown by the light-grey filled region. The extrapolation of the $\\gamma$-ray \\textit{Fermi} spectrum with its uncertainties is shown by the thin solid line with a hatched region. } \\label{fig:spectrum} \\end{figure}  From three statistically acceptable X-ray models, BB+PL, NSA+PL, and NSMAX+PL, considered in the previous Section, two latter result in similar parameters of the pulsar thermal emission. However, the NSMAX model is more justified from the physical reasons. Therefore in what follows we omit the NSA+PL model for simplicity.  In order to compare the X-ray data and the optical upper limit, the latter must be corrected for the interstellar extinction $A_V$. The standard $A_V$--$N_H$ relation \\citep{predehl1995AsAp} can be used to estimate $A_V$. The BB+PL and NSMAX+PL X-ray spectral fits suggest $N_H$ in a range of (1.0--2.3)~$\\times$ 10$^{21}$ cm$^{-2}$ (Table~\\ref{t:x-fit}). This corresponds to the $A_V$ range of 0.6$-$1.4. However, $A_V$ can hardly exceed the entire Galactic extinction in this direction of 0.8 recently estimated by \\citet{Schlaf}. Therefore, we accept 0.8 as a conservative extinction value for dereddening the optical upper limit. At the same time, the actual $N_H$ value can be larger than 1.4~$\\times$ 10$^{21}$ cm$^{-2}$ which corresponds to this $A_V$. For instance, \\citet{XraySecondDeLuca} estimate entire Galactic $N_H=(2.1\\pm0.2)\\times 10^{21}$~cm$^{-2}$ from the spectral analysis of extra-galactic sources in the pulsar field.  Considering the half-thickness of $\\sim$~100 pc for the Galactic gaseous disk responsible for the extinction, the pulsar latitude $b$~= $-$16$^{\\circ}$, and the minimal $A_V=0.6$ ($N_H=10^{21}$~cm$^{-2}$) we obtain the minimum distance to the pulsar of $\\sim$~270 pc. This value is derived assuming the uniform $A_V$ scaling with distance within the disk. The upper limit on the distance of $\\sim$900 pc was estimated by \\citet{XraySecondDeLuca} based on an assumption that the pulsar intrinsic $\\gamma$-ray luminosity cannot exceed its spin-down energy loss rate. The 500 pc distance accepted in Sect. 3.2 is consistent with these limits.  At the distance of 500 pc the optical (in the $V$ band) and nonthermal X-ray (in the range of 2--10 keV) luminosities of the pulsar are $L_{V}$~$\\la$ 1.1~$\\times$~10$^{27}$ erg~s$^{-1}$ and $L_{X}$~= 6.0~$\\times$~10$^{29}$ erg~s$^{-1}$. Accounting for the spin-down luminosity  $\\dot E= 5.8 \\times 10^{33}$ erg~s$^{-1}$ \\citep{BlindFreqAbdo2009} they yield  the optical and X-ray efficiencies of the pulsar $L_{V}/\\dot E$~$\\la$ 10$^{-6.7}$ and $L_{X}/\\dot E$~$\\approx$ 10$^{-4.0}$. These values are compatible with the empirical X-ray luminosity and efficiency {\\sl vs} age dependencies demonstrated by the pulsars detected in the optical and X-rays \\citep[e.g.,][]{zharikov2006AdSpR, zhar2}. This also supports the distance estimate of 500 pc.  In Fig.~\\ref{fig:spectrum} we compare the unabsorbed X-ray and $\\gamma$-ray spectra of the pulsar with the optical upper limit of 0.052 $\\mu$Jy derived from the GTC observations and dereddened with $A_V$~= 0.8. \\footnotetext{The data were unfolded and unabsorbed by applying the factor (unfolded unabsorbed model)/(folded absorbed model) in each spectral bin, assuming the best-fit BB+PL model. This procedure is analogous to how an unfolded spectrum is plotted by the \\texttt{Xspec} \\texttt{plot ufspec} command. Note that the data points obtained this way are model-dependent and for different models will follow the respective best-fit lines. } The optical upper limit is two orders of magnitude lower than it would be expected from the extrapolation of the PL spectral component to the optical range. According to Sect.~\\ref{s:x-ray} the PL component is essential to describe the high-energy tail of the pulsar X-ray spectrum. This suggests a spectral break in the PL component  between the optical and X-rays, as it is observed for all middle-aged pulsars detected in both domains \\citep{shibanov2006AsAp}. The extrapolation of the $\\gamma$-ray PL spectrum of the pulsar (see Fig.~\\ref{fig:spectrum}) lies well below the optical upper limit.  The best-fit NSMAX spectral components with variable and fixed normalisations are shown by thin and thick dashed lines in Fig.~\\ref{fig:spectrum}, respectively. For completeness we also show the NSA spectral components (dash-dotted lines). As seen, to significantly constrain the NS thermal emission in these models one has to go as deep as $\\sim$~30\\fm0 in the optical, which is not feasible with current instrumentation. The BB hot spot spectral component derived from the X-ray data with the BB+PL model (dotted line in the Fig.~\\ref{fig:spectrum}) all the more cannot be currently reached neither in the optical nor in UV. The light-grey region in Fig.~\\ref{fig:spectrum} contains all possibilities for the soft thermal component in the BB+BB+PL model for allowed $R^{\\infty}$ range in accordance with the 99\\% confidence contour of Fig.~\\ref{fig:weiss}. We may conclude that the optical upper limit does not put any additional constraints on thermal emission from the NS surface.  However, according to  Fig.~\\ref{fig:spectrum}, the entire surface thermal spectral component can be reached in UV. It can also dominate over the PL component there, if the PL component has approximately flat spectral slope from the optical to the UV, as it is observed for other middle-aged pulsars. The latter would be better constrained at longer optical wavelengths, less affected by the interstellar extinction. Therefore, UV observations of J0357$+$3205 would be useful to constrain its surface temperature. There are only few pulsars with thermal emission detected in the UV range namely PSR B0656+14 \\citep{Durant}, PSR B1055$-$52 \\citep{Mignani1055}, PSR J0437$-$4715 \\citep{Karg2004}, and Geminga \\citep{kargaltsev2005ApJ}. In addition, \\citet{Kaplan} reported detection of UV thermal emission from a few isolated neutron stars. In all these cases the UV data on thermal emission were of a great complement to the X-ray data.  Accounting for the direction of the pulsar proper motion and the spindown age, we find its likely birth place in the $\\lambda$-Orionis cluster, a 5 Myr active star forming region located in $\\sim$~32\\degs\\ from the pulsar and in $\\sim$~450$\\pm$50 pc from Earth \\citep{mayne2008MNRAS}. Several authors proposed that an expanding molecular ring surrounding the cluster is a supernova remnant left by a Type II supernova explosion of a massive companion of the O-type $\\lambda$~Ori star about 1 Myr ago \\citep[see e.g.,][]{cunha1996AsAp,dolan2002AJ}. Adopting this birth place we independently constrained the pulsar age of 0.2--1.3 Myr, accounting for the pulsar proper motion uncertainties, and the cluster and the pulsar distance ranges. This is consistent with its spindown age of 0.54 Myr.   \\label{sec4} \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.61]{coolcurves7.eps} \\end{center} \\caption{ Measured entire surface temperatures $T^{\\infty}$ as seen by a distant observer (filled circles) for NSs of different age $t$ in comparison with the cooling theory predictions (hatched regions). Dense hatched region correspond to the standard cooling theory, while sparse hatched region shows the minimal cooling theory predictions. The bold star with error-bars show the PSR~J0357$+$3205 surface temperature. } \\label{fig:cool} \\end{figure}  Accepting this age range and the NS effective temperature of 36$^{+8}_{-6}$ eV derived from the NSMAX+PL fit we can compare these with the NS cooling theory predictions. The J0357$+$3205 position on the temperature--age plane is shown in Fig.~\\ref{fig:cool} with the bold star with error-bars. The data for other isolated neutron stars (filled circles) are taken from \\citet{shternin2011MNRAS}. It is seen, that J0357$+$3205 is among the coldest cooling NSs known. The dense hatched region shows the range of the NS temperatures that can be obtained by the standard cooling theory where the modified Urca processes are considered as the main neutrino emission mechanism \\citep[e.g.,][]{yakovlev2004ARA}. The J0357$+$3205 position agrees well with the standard cooling theory. However, as seen from Fig.~\\ref{fig:cool} the standard cooling theory is insufficient to reproduce the data on all cooling NSs. Therefore, with sparse hatched region we show the range of cooling curves obtained within the minimal cooling scenario \\citep{gusakov2004AsAp,page2004ApJS,page2009ApJ} which takes into account the presence of the baryon superfluidity inside neutron stars. In this scenario the specific process of the neutrino emission due to a Cooper pair formation cools the star more effectively than the modified Urca process. To date, the parameters of the superfluidity can be plausibly adjusted  (sparse hatched region) to fit all the data on the observed NSs temperatures \\citep{gusakov2004AsAp}, including the likely rapidly cooling NS in Cas A \\citep{shternin2011MNRAS}. Obviously, J0357$+$3205 agrees with the minimal cooling scenario as well.  At the same time, according to Fig.~\\ref{fig:weiss}, the entire surface temperature in the blackbody spectral model is poorly constrained, taking into account the uncertainties in the NS radius and distance to the pulsar. Thus it is not possible to extract any valuable information from comparison of the BB+BB+PL fit results with the cooling theories.  To summarise, our deep optical observations of PSR J0357$+$3205 allowed us to constrain  the pulsar nonthermal emission, suggesting a strong spectral break in this emission between the optical and X-rays. Reanalysis of X-ray data allowed us to constrain the NS thermal spectrum and to measure the effective temperature of the NS surface $T^{\\infty}$~= 36$^{+8}_{-6}$ eV. Comparing the optical upper limit with the NS thermal spectrum we conclude that the thermal emission from the entire surface of the NS can be feasibly examined in the UV range and likely dominate there over the nonthermal emission and the emission from pulsar hot spot(s). This makes J0357$+$3205 a promising target for UV observations."
    },
    "1402/1402.5699_arXiv.txt": {
        "abstract": "Unlike general relativity, the scalar gravitational waves can be excited due to the radial oscillations in scalar-tensor gravity. To examine the scalar gravitational waves in scalar-tensor gravity, we derive the evolution equations of the radial oscillations of neutron stars and determine the specific oscillation frequencies of the matter oscillations and scalar gravitational waves, where we adopt two different numerical approaches, i.e., the mode analysis and direct time evolution. As a result, we observe the spontaneous scalarization even in the radial oscillations. Depending on the background scalar field and coupling constant, the total energy radiated by the scalar gravitational waves dramatically changes, where the specific oscillation frequencies are completely same as the matter oscillations. That is, via the direct observations of scalar gravitational waves, one can not only reveal the gravitational theory, but also extract the radial oscillations of neutron stars. ",
        "introduction": "\\label{sec:I}  Since general relativity has been proposed, many experiments have been performed to verify the gravitational theory. Most of these attempts are done in the weak gravitational field such as our solar system, but nothing indicates the failure of general relativity. On the other hand, since the astronomical observations in the strong gravitational field are very poor, the gravitational theory in such a strong-field regime could be still unconstrained. That is, the gravitational theory to describe the phenomena in the strong-field regime might be different from general relativity, and one might be possible to probe the gravitational theory via the observations of the deviation from general relativity. In practice, up to now there are many suggestions to observationally test the gravitational theory in the strong-field regime \\cite{W1993,W2001,P2008}. The technology is developing more and more, which will enable us to accurately observe the phenomena in the strong-field regime. These coming new observations might be possible to use as the test of gravitational theory.   So far, a lot of alternative gravitational theories are proposed. Among them, scalar-tensor theory is the one of the simplest alternative gravitational theories, which must be a natural extension of standard general relativity \\cite{DE1992}. One of the motivations to consider scalar-tensor gravity is that this theory can be obtained in the low energy limit of string and/or other gauge theories. In scalar-tensor gravity, the scalar field plays an essential role in addition to the usual tensor field in general relativity, where the matter field is described by using the effective metric $\\tilde{g}_{\\mu\\nu}$ associated with the scalar and gravitational field, $\\varphi$ and $g_{*\\mu\\nu}$, via the conformal transformation, i.e., $\\tilde{g}_{\\mu\\nu}=A^2(\\varphi)g_{*\\mu\\nu}$. The Brans-Dicke theory \\cite{BD} is the simplest version of scalar-tensor gravity, where $A(\\varphi)$ is defined as $A(\\varphi)=exp(\\alpha\\varphi)$. The coupling parameter $\\alpha$ can be associated with the Brans-Dicke parameter $\\omega_{\\rm BD}$ as $\\alpha^2=1/(2\\omega_{\\rm BD}+3)$, which are constrained through the solar system experiments, i.e., $\\omega_{\\rm BD} \\gsim 40000$ which is corresponding to $\\alpha < 10^{-5}$ \\cite{EF2004}. Within this restriction on the coupling parameter, it is almost impossible to predict a large deviation from general relativity in the strong-field regime.   A different functional form of the conformal factor is also suggested by Damour and Esposito-Far\\`{e}se \\cite{DE1993,DE1996}, where $A(\\varphi)\\equiv exp(\\alpha\\varphi + \\beta\\varphi^2/2)$. With this type of coupling, even if $\\alpha$ is almost zero, the relativistic stellar models in scalar-tensor gravity can significantly deviate from the predictions in general relativity. Additionally, they found that the stellar models in scalar-tensor gravity suddenly deviate from those in general relativity for the specific values of coupling parameters, which is referred to as {\\it spontaneous scalarization}. With respect to this phenomenon, Harada systematically examined with the technique of catastrophe theory and found that the spontaneous scalarization can happen for $\\beta\\lsim -4.35$ \\cite{H1998}. Recently, it is found that the spontaneous scalarization are possible for larger value of $\\beta$ in fast rotating relativistic stars \\cite{Daniela2013} and in the neutron star binary system \\cite{BPPL2013,Shibata2013,PBPL2013}. On the other hand, using the observations of pulsar white dwarf binary, Freire {\\it et al}. set a severe constraint on $\\beta$, i.e., $\\beta\\gsim -5$ \\cite{F2012}. Additionally, it is reported that $\\beta$ could be constrained to be larger than $-4.5$, depending on the equation of state \\cite{Shibata2013}. Maybe, although the constraint on $\\beta$ would become severer via the future observations, we focus on the range of $\\beta \\gsim -5$ in this paper.    The several attempts to observationally distinguish scalar-tensor gravity from general relativity have been already done in the past by using the redshift in the absorption lines of the X and $\\gamma$ rays emitted from the stellar surface \\cite{DP2003}, the spectrum of the gravitational waves radiated from relativistic stars \\cite{SK2004,SK2005}, and the rotational effect around compact objects \\cite{Sotani2012}. In this paper, we consider the different approach, i.e., the scalar gravitational waves driven by the radial oscillations. In fact, the gravitational waves can not be excited due to the radial oscillations in general relativity. This means that the detection of scalar gravitational waves itself becomes the proof the existence of scalar field. From the observational point of view, if the scalar gravitational waves exist, one could in principle identify the scalar gravitational waves with more than three gravitational wave detectors, because we have only three degrees of polarizations in scalar-tensor gravity, such as two usual tensor gravitational waves and scalar gravitational wave. In fact, we expect that five gravitational wave detectors will be in operation in the future, such as two advanced LIGOs \\cite{aLIGO}, advanced Virgo \\cite{aVirgo}, KAGRA \\cite{KAGRA}, and IndIGO \\cite{IndIGO}. On the other hand, the method how to separate and reconstruct an arbitrary number of polarization modes by using the observational data by multiple interferometric gravitational wave detectors, is also developing, which is a model-independent approach \\cite{HN2013}. It should be notices that the scalar gravitational waves in physical frame are proportional to the cosmological value of scalar field \\cite{BPPL2013,Shibata2013}, whose value must be quite small. That is, if the scalar gravitational waves exist, they might be quite weak and difficult to detect in the current gravitational wave detectors.   The radial oscillations of relativistic stars in general relativity have been examined since early times \\cite{C1964,GL1983,GL1992,KR2001} in the context of the stability analysis. Meanwhile, the scalar gravitational waves in scalar-tensor gravity are also examined in the black hole formation due to the dust collapse \\cite{Shibata1994,Scheel1995,H1997} and the test particle around a Kerr black hole \\cite{Saijo1997}. Anyway, this is the first time to calculate the scalar gravitational waves driven by the stellar radial oscillations in scalar-tensor gravity suggested by Damour and Esposito-Far\\`{e}se \\cite{DE1993,DE1996}. For this purpose, we will derive the perturbation equations of radial oscillations and make numerical calculations to examine it.    This paper is organized as follows. In the next section, we briefly mention the equilibrium of nonrotating relativistic stars in scalar-tensor gravity. In Sec. \\ref{sec:III}, we derive the perturbation equations describing the radial oscillations of relativistic stars in scalar-tensor gravity. The numerical results are shown in Sec. \\ref{sec:IV}, where the specific frequencies are determined with the mode analysis and the direct time evolution. Then, we make a conclusion in Sec. \\ref{sec:V}. We adopt the geometric units, $c=G_*=1$, where $c$ and $G_*$ denote the speed of light and the gravitational constant, respectively, and use the metric signature is $(-,+,+,+)$. ",
        "conclusions": "\\label{sec:V}  Neutron stars are one of the best candidates to probe the gravitational theory in the strong-field regime. In this paper, we especially focus on the radial oscillations of neutron stars in scalar-tensor gravity to examine the scalar gravitational waves driven by the matter oscillations. In fact, the gravitational waves are not excited due to the radial stellar oscillations in general relativity, while one can expect to observe the scalar gravitational waves due to such oscillations in scalar-tensor gravity. For the calculations of the scalar gravitational waves, we first derive the perturbation equations for radial oscillations in scalar-tensor gravity. From the equation of system of radial oscillations, we find that the matter oscillations depend only on the background scalar field, independently of the scalar gravitational waves. On the other hand, the wave equation of scalar gravitational waves has a source term composed of the matter oscillations. Due to such a specific coupling, we can determine the frequencies of matter radial oscillations by the mode analysis. As a result, we show that the spontaneous scalarization can be observed even in the radial oscillations, which might enable us to find the imprint of gravitational theory with the help of the other observations such as stellar mass and/or compactness.    Additionally, to examine the scalar gravitational waves driven by the matter radial oscillations, we directly make a numerical simulation of the evolution equations, where we fix the initial energy of matter oscillations to be $10^{-4}M_\\odot$.  Then, we find that the scalar gravitational waves can be excited if the background scalar field exists. We also find that the total energy radiated by the scalar gravitational waves depends strongly on the background scalar field and the coupling constant $\\beta$, where the massive star has a potential to radiate more energy of scalar gravitational waves. Furthermore, we make the fast Fourier transform to see the specific oscillation frequencies of radiated scalar gravitational waves, which are exactly same as the frequencies of matter oscillations. That is, via the observations of scalar gravitational waves, one can extract the frequencies of stellar radial oscillations. This is an advantage in scalar-tensor gravity, because the radial gravitational waves can not be excited in general relativity as mentioned before. We remark that one might have another chance to observe an imprint of radial oscillation in the gravitational waves, if radial oscillations are strongly excited, for example in core-collapse supernovae, and those oscillations make nonlinear coupling with nonradial oscillations, where the oscillations with combination frequencies could be excited \\cite{PSN2007}. If so, one might be possible to measure the background scalar field via such nonlinear coupling. In this paper, as a first step, we neglect the effects of the solid crust layer, magnetic fields, and the exotic matter inside the star, which are also important properties of neutron stars. Such effects could bring us the additional information about the stellar properties \\cite{SKLS2013,Sotani2007,Sotani2011,Yasutake}, which might make observational constraints in the gravitational theory stronger."
    },
    "1402/1402.3240_arXiv.txt": {
        "abstract": "FU Orionis stars (FUORS) are eruptive pre-main sequence objects thought to represent quasi-periodic or recurring stages of enhanced accretion during the low-mass star-forming process. We characterize the sample of known and candidate FUORS in an homogeneous and consistent way, deriving stellar and circumstellar parameters for each object. We emphasize the analysis in those parameters that are supposed to vary during the FUORS stage. We modeled the SEDs of 24 of the 26 currently known FUORS, using the radiative transfer code of \\citet{whitney2003b}. We compare our models with those obtained by \\citet{robitaille2007} for Taurus class~II and I sources in quiescence periods, by calculating the cumulative distribution of the different parameters. FUORS have more massive disks: we find that $\\sim80\\%$ of the disks in FUORS are more massive than any Taurus class~II and I sources in the sample.  Median values for the disk mass accretion rates are  $\\sim$ 10$^{-7}$ $\\msun$/yr vs $\\sim$ 10$^{-5}$ $\\msun$/yr for standard YSOs (young stellar objects) and FUORS, respectively. While the distributions of envelope mass accretion rates for class I FUORS and for standard class I objects are similar, FUORS, on average, have higher envelope mass accretion rates than standard class~II and class~I sources. Most FUORS ($\\sim$ 70\\%) have envelope mass accretion rates above $10^{-7}\\,\\msun$/yr. In contrast, 60\\% of the classical YSO sample have accretion rates below this value. Our results support the current scenario in which changes experimented by the circumstellar disk explain the observed properties of these stars. However, the increase in the disk mass accretion rate is smaller than theoretically predicted \\citep{frank1992,hartmann1996a}, though in good agreement with previous determinations. ",
        "introduction": "FU Orionis stars (FUORS) are a class of variable young stellar objects that show brightness variations of the eruptive type \\citep{herbig1977}. The main feature observed in these variables is a sudden increase in brightness  ($ 3-6\\,$mag in the optical), in an elapse of time of a few months. This episode is known as the ``outburst'', after which the object remains bright for years or decades, and then fades in a few centuries back to the pre-outburst stage. The outburst, however, occurs in a different way for each FUORS (see, e.g., \\citealt{hartmann1996a,clarke2005a}).  These stars exhibit several indicators of youth, such as the presence of the lithium $ 6707\\,$\\AA~ line in optical spectra, and the association with reflection nebulae and infrared excesses originating from dust grains in circumstellar disks. Moreover, they are spatially and kinematically related to known star-forming regions and in some cases, FUORS have high extinction values in the optical suggesting that they are still embedded in the parent cloud material (see, e.g., \\citealt{hartmann1996a}).  FUORS show several properties that strongly suggest the presence of a circumstellar disk, such as broad spectral energy distributions (SEDs, \\citealt{kenyon1988b}), stellar spectral types that become progressively colder with increasing wavelength \\citep{hartmann1985,kenyon1988b}, spectral linewidths that increase with decreasing wavelength \\citep{hartmann1987a,hartmann1987b}, double-peaked line profiles in high-resolution optical and near-infrared (NIR) spectra \\citep{hartmann1985,kenyon1988b} as well as P-Cygni profiles with no evidence for redshifted emission or absorption \\citep{kenyon1988b,hartmann1995}, and finally, deep, broadened infrared CO bands in absorption \\citep{kenyon1988a,reipurth1997b}.  Another class of eruptive variables are the so-called EXOR stars, named after EX Lup, the progenitor of the class \\citep{herbig1989,herbig2008}. Their optical brightness increases by $ 1-4 $\\,mag on time scales of weeks or months, then fading back during a few months to its original state after. During its low activity stage they exhibit T~Tauri-like characteristics, while during the outburst stage they usually display emission in the optical spectrum as well as in the infrared CO bandheads (e.g., \\citealt{aspin2010}).  To reproduce the SEDs of FUORS, modelers use dusty disks and infalling envelopes (e.g., \\citealt{hartmann1985,kenyon1988b,calvet1991,hartmann1995,calvet1998,whitney2003a,whitney2003b}). Indeed the presence of a circumstellar disk is essential to explain the FU Orionis phenomenon. The disk is where material coming from the surrounding infalling envelope accumulates, heats-up, and finally destabilizes the structure of the disk itself, causing a thermal \\citep{frank1992, bell1994}  and/or a gravitational \\citep{zhu2009,zhu2010,vorobyov2005,vorobyov2006,vorobyov2010} instability that eventually leads to the characteristic outburst.  During this episode an increase of the brightness takes place, affecting mainly the optical wavelengths since the excess emission comes from the inner regions of the disk, which are heated by the viscous dissipation released after the instability has triggered an increase in the disk accretion rate. \\citet{frank1992} suggests that the central objects of the FUORS systems alternate between low ($10^{-7}\\,$\\msun/yr) and high ($10^{-4}\\,$\\msun/yr) mass accretion rates. The former corresponds to a low activity, quiescent state, while the latter corresponds to periods of high activity. Furthermore, the transformations undergone by the disk are what cause  the observed phenomenon. Similarly, EXOR events are also attributed to thermal instabilities in the inner disks \\citep{aspin2011a}.  Currently, 26 FUORS have been identified and classified as class~I or class~II objects according to the shape of their SEDs \\citep{lada1987}. This includes the ``confirmed'' FUORS, for which the sudden increase in brightness has been recorded, and the ``candidate'' FUORS, which share many, if not all, of the properties of bona-fide FUORS but for which an outburst has not been observed. In this paper we model and analyze the SEDs of 24 confirmed and candidate FUORS to determine the physical and geometrical parameters of the star and the disk.  In Section 2 we present the sample, describe the adopted model and outline the procedure used in the SED modeling. In Section 3 we analyze the individual sources. Our results are presented in Section 4. Finally in Section 5 we summarize the results and conclusions. ",
        "conclusions": "\\label{s:summary}  In this work we present the modeling of the SEDs of a sample of 24 class~II and class~I FU Orionis stars. These SEDs were constructed from fluxes obtained from the literature (Table~\\ref{t:fluxes}), including Spitzer-IRS infrared spectra in the $5-35\\,\\mu$m range for V1515~Cyg, BBW~76, FU~Ori, V346~Nor and V1057~Cyg, and in the $5 -14\\,\\mu$m range for RNO~1B, RNO~1C, L1551~IRS5 and Par~21 \\citep{green2006,quanz2007b}. For Re~50~N~IRS1, we used an ISO-SWS spectrum in the $5 -15\\,\\mu$m range obtained by \\citet{quanz2007b}.  Initially we modeled each source applying the grid of \\citet{robitaille2006}, to later use these models as starting points for a more refined analysis using the code of \\citet{whitney2003b}. The parameters corresponding to the best model fits are given in Tables \\ref{t:resclass2} and \\ref{t:resclass1} for class~II and class~I FUORS, respectively. Figures \\ref{f:v1515cyg} to \\ref{f:v2775} show the corresponding SEDs. For sources V1515~Cyg, BBW~76, PP~13S, V1647~Ori, FU~Ori, V1057~Cyg, Z~CMa, L1551~IRS5, ISO-Cha~192, V2492~Cyg, V1331~Cyg, Par~21, and V2775~Ori we compared our parameters values with those derived by other authors, finding in general a good agreement. For the remaining 11 sources, this is the first time a model of their SED is derived.  Figure \\ref{f:cumul} shows the accumulative distribution functions of disk masses, disk accretion rates, envelope accretion rates and stellar temperatures of FUORS in our sample and standard class~II and class~I objects in a quiescence state from \\citet{robitaille2007}. Table \\ref{t:median} gives the median values for both groups. The comparison shows that:  \\begin{enumerate}   \\item On average FUORS disks are more massive than standard class~II and class~I objects disks. About 80\\% of FUORS disks have masses $> 0.1\\,\\msun$, while standard class~II and class~I objects have disk masses $< 0.06\\,\\msun$. \\item Disks mass accretion rates are higher for FUORS than for classical YSOS. The great majority of FUORS ($\\sim90\\%$) have $\\dot{\\rm M}_{disk}>10^{-6}$\\,M$_{\\odot}$/yr, while $\\sim95\\%$ of the standard class~II and class~I objects have $\\dot{\\rm M}_{disk}<10^{-6}$\\,M$_{\\odot}$/yr. Median disks accretion rates are $\\sim 10^{-5}$\\,M$_{\\odot}$/yr vs $\\sim 10^{-7}$\\,M$_{\\odot}$/yr for FUORS and classical YSOs, respectively. \\item The distributions of envelope accretion rates for class I FUORS and standard class I objects are indistinguisable. Most FUORS ($\\sim70\\%$) have envelope accretion rates $> 10^{-7}\\,\\msun$/yr. Median envelope accretion rates are $\\sim 10^{-6}$\\,M$_{\\odot}$/yr vs $\\sim 10^{-8}$\\,M$_{\\odot}$/yr for FUORS and standard YSOs, respectively. \\item The distribution of stellar temperatures for FUORS and classical YSOs are similar in shape, but the FUORS are shifted $\\sim2000\\,$K to higher temperatures.  \\end{enumerate}  The cumulative distributions for confirmed and candidate FUORS (see Table~\\ref{t:sample}) show no significant differences, suggesting that most candidate objects, in fact, belong to the FUORS class. We caution, however, on the small number of objects in each class (14 confirmed and 10 candidate FUORS).  For the seven objects in our sample, for which we have SEDs both in the outburst and in the quiescence stage (see Tables~\\ref{t:resclass2} and \\ref{t:resclass1}), 2 class II (RNO 1B, and V1647 Ori), and 5 class I (OO Ser, V1313 Cyg, V2492 Cyg, HBC 722, and V2775 Ori), we note that while the disk and stellar parameters show variations, the envelope parameters ($\\dot{M}$, R$ _{\\rm max}$, and M$_{env}$) do not change, suggesting the outbursts are triggered by an instability after a long build-up phase.  The current scenario of FUORS events states that the circumstellar disk of a YSO builds up material injected from the envelope until it becomes thermally \\citep{frank1992, bell1994, hartmann1996a} and/or gravitationally \\citep{zhu2009,zhu2010,vorobyov2005,vorobyov2006,vorobyov2010} unstable. In particular, using the model parameters and disk properties listed in Tables~\\ref{t:resclass2} and \\ref{t:resclass1} we calculated the Toomre $ Q $ gravitational stability parameter \\citep{Toomre1964}.  If $Q < 1$, the disk is unstable. Most of the models are unstable for $R > 5-20$\\,AU. The only exception is L~1551~IRS~5, which is gravitationally unstable at a larger scale ($R \\gtrsim 50$\\,AU). Nevertheless all disk are unstable well inside the centrifugal radius (see Table \\ref{t:averages}). Consequently, gravitational instabilities may contribute to the outburst eruptions, in addition to thermal instabilities, resulting in an increase of the mass accretion onto the central object.  What we have described so far agrees with this picture. However, the disk mass accretion rate $\\dot{\\rm M}_{disk}\\sim 10^{-5}\\,\\msun$/yr we obtain (see Table \\ref{t:median}) is one order of magnitude lower than the $\\sim 10^{-4}\\,\\msun$/yr predicted by the theory \\citep{frank1992,hartmann1996a}. Nevertheless, previous models of individual FUORS objects obtain $\\dot{\\rm M}_{disk}$ values consistent with those presented in this work (see, e.g., \\citealt{pfalzner2008, aspin2008}).  Although the average values for the parameters for both groups of FUORS are similar to those theoretically expected, the individual values listed for each object in Tables \\ref{t:resclass2} and \\ref{t:resclass1} differ significantly. This can be in part attributed to the fact that the group of the FU Orionis stars itself is not an homogeneous sample. While they all share a particular set of characteristics, those appear in different ways for each object. For instance, while all FUORS show a sudden brightness increase of several magnitudes, followed by a slow decrease to their previous state, the way the brightness jump develops in time is different for each object. A clear example of this is the great diversity in the light curves of the three prototypes of the class \\citep{hartmann1996a}. It is therefore reasonable to expect that the values of individual parameters of each member of the group will ultimately differ.  Lastly, we would like to draw attention to three sources in particular. V1647~Ori is a special case on the FUOR sample since it has been well studied before and after the outburst, having SEDs for both epochs. This makes that source a prime candidate for the study of the FU Orionis event, though it has to be approached with care, since its classification as a FUOR or EXOR is still under debate \\citep{aspin2006,aspin2011b,semkov2012a}. The other two particular sources are V2492~Cyg and HBC~722. These objects are, at the moment of writing, the last two for which an FUORS-like outburst has been observed. They show characteristics proper of bona-fide FUORS, as shown in Sections~\\ref{sec-v2492cyg} and \\ref{sec:hbc722}, however its inclusion in the FUORS class is still not certain. Nevertheless, our SED modeling shows behaviors similar to V1647~Ori, the former newest member of the class.  V1647~Ori and V2492~Cyg also show EXOR characteristics, and from our modeling we see that they do not show a large variation between the outburst and the quiescent phases. For example, their disk masses do not change with the outburst (see Tables~\\ref{t:resclass2} and \\ref{t:resclass1}), and the variation of the disk mass accretion rate is lower than for other FUORS. Nevertheless, when compared with other FUORS, the parameter values derived for those sources are still within the range established by the rest of the FUORS sample, and could then be considered FUORS. However, if we had just analyzed only those two sources while taking into consideration that EXOR outbursts are thought to be ``scaled-down'' versions of FUORS outbursts, it is very likely that they would have been considered EXORS. This shows the uncertainty and difficulty of disentangle the two types of outburst episodes.  Despite sharing common properties, each FUORS or FUORS candidate has its own peculiarities that are not currently well understood. It is therefore of great interest to study the most extreme objects of the class to reach a full understanding of this period of great activity in circumstellar disks.  The work we have presented here is the first compilation of SEDs of the currently known FUORS. Of the 26 currently known FUORS, 2 do not have enough observations as to construct the SEDs, and thus are not analyzed. For 21 of the remaining 24 we compile in one place the observations taken in all wavelengths, producing the most complete SEDs possible so far. For 3 FUORS (AR 6A, AR 6B, V2492 Cyg), SEDs for only a limited range ($\\lambda$ $<$ 20 $\\mu$m) were constructed and thus values for the derived parameters are not fully determined.  Finally, for 11 of the 24 FUORS analyzed (V1735 Cyg, V883 Ori, RNO 1B, RNO 1C, AR 6A, AR 6B, V900~Mon,  V346~Nor, OO~Ser, RE 50 N IRS 1 and HBC~722) we provide for the first time a complete SED modeling to determine the physical and geometrical parameters of the star$+$disk$+$envelope system. Furthermore, this is the first time all the known FUORS with an observed SED are modeled with the same code at the same time, providing an homogeneous set of results. The data we present here will be of great help for future studies in the field."
    },
    "1402/1402.3595_arXiv.txt": {
        "abstract": "In order to empirically determine the timescale and environmental dependence of stellar cluster disruption, we have undertaken an analysis of the unprecedented multi-pointing (seven), multi-wavelength (U, B, V, H$\\alpha$, and I) Hubble Space Telescope imaging survey of the nearby, face-on spiral galaxy M83.  The images are used to locate stellar clusters and stellar associations throughout the galaxy. Estimation of cluster properties (age, mass, and extinction) was done through a comparison of their spectral energy distributions with simple stellar population models. We constructed the largest catalog of stellar clusters and associations in this galaxy to-date, with $\\sim1800$ sources with masses above $\\sim5000$~\\msun and ages younger than $\\sim300$ Myr. In the present letter, we focus on the age distribution of the resulting clusters and associations.  In particular, we explicitly test whether the age distributions are related with the ambient environment. Our results are in excellent agreement with previous studies of age distributions in the centre of the galaxy, which gives us confidence to expand out to search for similarities or differences in the other fields which sample different environments. We find that the age distribution of the clusters inside M83 varies strongly as a function of position within the galaxy, indicating a strong correlation with the galactic environment.  If the age distributions are approximated as a power-law of the form $\\dndt \\propto t^{\\zeta}$, we find $\\zeta$ values between $0$ and $-0.62$ ($\\zeta \\sim -0.40$ for the whole galaxy), in good agreement with previous results and theoretical predictions. ",
        "introduction": "\\label{sec:intro}  The age distribution of stellar cluster populations in galaxies is the result of the combined effects of the cluster formation history and cluster disruption. Studies have shown a tight correlation between the maximum intensity of the star formation history and the age distribution of the cluster population in multiple galaxies (e.g. NGC~7252 - \\citet{miller97}; \\citet{schweizes98}, \\citet{chien10}; M82 - \\citet{konstantopoulos09}).  More quiescent galaxies, such as the Milky Way (at least in the solar neighbourhood) have a smoother, rather flat, cluster age distribution \\citep[e.g.][]{lamers05,piskunov06}.  The age distribution of a star cluster population (\\dndt) is defined as the number of clusters observed within some linear time interval, and, at present, the role of cluster disruption in setting the shape of the age distribution is still an open topic of debate. If cluster disruption has any dependency with environment, we would expect to see the age distribution changing over different locations, e.g. where there are differences in gas density, or in the tidal field. However, if cluster disruption is independent of environment, the age distributions of any location should be similar.  We chose the galaxy M83 because it presents such differences in environment, allowing us to test these assumptions \\citep[see e.g.][for differences in gas densities across M83]{lundgren04}.  Previous studies \\citep[e.g.][]{fall05} of the Antennae galaxies (a currently starbursting, merging galaxy) have found that the cluster age distribution is quite steep, with \\dndt$ \\sim t^{-1}$, and this has been interpreted as evidence for rapid and strong cluster disruption under the assumption that the cluster formation rate/history has been constant over the past few hundred Myr \\citep{fall05,whitmore07}.  However, \\citet[][among others]{bastian09} showed that the steepness of the age distribution can be affected by the star formation rate. The age distribution becomes shallower when corrected for an increasing star formation rate, limiting the role of cluster disruption in shaping the age distribution (assuming that both are studied under the same time interval). Also, based on theoretical studies \\citep[e.g.][and references there in]{elmegreen10,kruijssen11} it was showed that the ambient environment where a cluster resides can drastically affect the steepness of the age distribution, making it disrupt faster when the surface gas density ($\\Sigma_{gas}$) is higher, and live longer when $\\Sigma_{gas}$ is low, indicating a strong relation with environment.  \\citet{chandar06} used the publicly available catalogue of clusters compiled by \\citet{hill06} to estimate the age distribution of clusters in the Small Magellanic Cloud (SMC). The authors used a lower mass limit of $10^3$\\msun\\ and ages between $7\\le Log(\\tau /yr) \\le 9$. They found a steep distribution with \\dndt$\\sim t^{-0.85}$ in the age interval mentioned.  However, using the same catalogue, \\citet{gieles07} and \\citet{degrijs08} showed that the steepness in the age distribution was caused by the sample of Hill \\& Zaritsky being luminosity-limited. When a mass cut above the completeness was used, the resulting distribution was flat, leading to the conclusion that cluster disruption has not significantly altered the age distribution of SMC clusters.  The Large Magellanic Cloud (LMC) hosts a much larger cluster population than the SMC, and has been the subject of numerous studies. \\citet{chandar10} used the \\citet{hunter03} catalogue to study the age distribution, and supplement this catalogue with objects that were avoided by Hunter et al., due to nebular emission.  They find a steep age distribution, \\dndt$\\sim t^{-0.8}$. However, \\citet{baumgardt13} and \\citet{degrijs13} combine all publicly available catalogues of clusters in the LMC and do not confirm the Chandar et al. findings. Instead they find much flatter results, again suggesting the cluster disruption has not had a strongly effect on the cluster population.  \\citet{silvavilla11} studied the age distributions of clusters in five near-by, face-on spiral galaxies, including M83 (NGC~5236) and NGC~1313. Fitting over the age range of 10 to 500~Myr, they found $\\dndt \\sim t^{-0.25}$ and $t^{-0.65}$ for M83 and NGC~1313, respectively. Their results suggest that the age distribution of star clusters might vary as a function of environment  The Panchromatic Hubble Andromeda Treasury survey \\citep[PHAT survey,][]{dalcanton12}, one of the largest studies to-date, has catalogued stellar clusters within M31 \\citep[][to date only $\\sim25$\\% of the survey area has been presented]{johnson12}. Fouesneau et al.~(2014) derived the ages and masses of the clusters in the survey and found a flat \\dndt\\ distribution for ages younger than $\\sim$70 Myr, again indicating the small role of disruption in setting the shape of the age distribution in this galaxy.  In the present work we study the age distributions of clusters in the nearby \\citep[$\\sim4.5$~Mpc,][]{thim03}, face-on spiral galaxy M83, based on seven pointings observed with the Wide Field Camera 3 (WFC3) onboard the Hubble Space Telescope (HST).  This study aims to complement and extend the previous studies of \\citet[][hereafter C10]{chandar10}, \\citet[][hereafter F12]{fouesneau12} and \\citet[][hereafter B11 and B12, respectively]{bastian11,bastian12}.  C10 reported a steep age distribution of clusters and associations in the first of the seven fields analyzed here (F1, see Fig.~1), with $\\dndt \\sim t^{-0.8}$.  B12 analyzed the same region (F1) and came to similar conclusions. However, they included another pointing further away from the galaxy centre (F2), and found $\\dndt \\sim t^{-0.5}$.  Here, we extend the B11 and B12 analysis by including five further WFC3 pointings of M83. For this study, we try to avoid the problem introduced by a luminosity-limited sample and we will only compare age distributions of mass-limited samples.  This paper is organized as follows.  In \\S~\\ref{sec:obs} we present the data and techniques used and in \\S~\\ref{sec:results} we show our main results.  We discuss the implications of our findings in \\S~\\ref{sec:discussion}.  \\begin{figure} \\includegraphics[width=8.5cm]{figure1.eps} \\caption{M83 R-band image \\citep{meurer06} with overlapping lines representing the areas covered with the HST. Taken from NASA/IPAC Extragalactic Database (NED).} \\label{fig:mosaic} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  \\begin{figure} \\includegraphics[width=8.5cm]{figure3.eps} \\caption{The indices from maximum likelihood fits to each field of the form $\\dndt \\sim t^{\\zeta}$.  The symbols represent different choices of the classes of sources used in the fits and different age ranges where the fit is made. Note that the fields are ordered in increasing $\\zeta$ values.  Numbers (N) at the top are the number of clusters for the filled circles in the legend.  ``Total\" represents the full catalogue.  Additionally, we plot the results from the C10 catalogue (with symbols that represent both class 1 and class 2 sources). The error is a combination of the standard deviation of the scatter and the random error in the fit.} \\label{fig:dndt2} \\end{figure}  As discussed in \\S~1, the age distribution of clusters is a convolution of the cluster formation history and cluster disruption. \\citet{fall05,whitmore07,chandar10,fall12}, among other authors, have suggested a Universal (or quasi-Universal) age distribution in a form of a steep power-law with an index $\\zeta$ between -0.8 and -1.0.  These authors found this in a variety of galaxies, therefore they conclude that this rapid disruption must be independent of the environment. Since it is unlikely that all galaxies would have undergone a rapid increase in the cluster formation rate over the past few hundred Myr, they argue that the steepness of the distributions is caused by rapid cluster disruption. However, if true, there are important discrepancies between this environmental independence and the results from analytic models based on N-body simulations \\citep[e.g.][]{baumgardt03,gieles06}, and it will be require a new framework to understand the evolution of star clusters.  It is worth noting that the Universal model was originally developed to explain the stellar cluster population of the merging, starburst galaxy, the Antennae, where it was assumed that the cluster formation rate and history was constant over the past few hundred Myr, contrary to expectations of numerical simulations \\citep[e.g.][]{baumgardt03,bastian09,chien10,kruijssen11}.  We have attempted to reduce systematics between the fields, by adopting $(i)$ the same photometric procedure, $(ii)$ the same SED models, and $(iii)$ fitting procedure using identical age bins. Additionally, we have also carried out a maximum likelihood fit to the distributions to avoid binning altogether. Our results from these two different methods are in excellent agreement.  The results found for the seven fields in M83 do not agree with the predictions of the Universal model, as we find significantly shallower age distributions, with values of $\\zeta$ between $-0.1$ and $-0.61$. These variations indicate a strong correlation with the local ambient environment.  These shallower distributions agree with analyses of the SMC \\citep{gieles07,degrijs08,portegieszwart10}, LMC \\citep{baumgardt13,degrijs13}, M51 \\citep{gieles05}, and NGC~2997 (Ryon et al. in prep.) cluster populations, as well as previous M83 results \\citep[B11; B12;][]{silvavilla11}.  Is there a way in which the observations presented here could be made to fit with the Universal scenario?  One could argue that the star formation history in M83 has been decreasing over the past 200-300~Myr, which would counteract the observational effects of cluster disruption.  In order to change the observed \\dndt\\ distributions to be consistent with the Universal expectations, the star formation rates would have had to have decreased by a factor of 2-3 (Field~1) to 10 (Fields 3 and 7) over the past 100~Myr. However, based on the results presented by \\citet{silvavilla11} and B12, based on resolved stellar populations, such changes are unlikely.  Under the assumption that the cluster formation history/rate has not dramatically changed in the time interval we are studying here \\citep[based on resolved stellar population studies, e.g.][]{silvavilla11}, we conclude that the cluster population of M83 has not been dramatically affected by cluster disruption, in contrast to expectations from the Universal scenario, where clusters are disrupted at a high rate independent of their ambient environment \\citep[e.g.][]{chandar06,chandar10,fall12}.  On the other hand, the mass and environmentally dependent disruption model \\citep[MDD model,][]{lamers05} predicts that the slope of the age distribution will be strongly tied to the tidal fields and the surface density of giant molecular clouds (GMCs), leading to a strong correlation with the ambient environment. Under the assumption of a (roughly) constant cluster formation history, if the tidal field is strong and the GMC density is high, cluster disruption is expected to be stronger, leading to steeper age distributions. On the contrary, if the tidal field is weak and the GMC density is low, there will be less disruption, and the age distribution will be flatter. We find that the inner fields, where the tidal forces and GMC density are highest \\citep[c.f.][]{bastian12} have steeper age distributions, while in the outer regions they are flatter, in excellent agreement with predictions from the models \\citep[e.g.][]{lamers05,kruijssen11}.   We will further quantify the relation between environment and disruption using the M83 cluster dataset in a future study.  A study of the resolved stellar population, along with the presentation of the full catalogue and a further study of cluster disruption, will be presented in a future work, Silva-Villa et al. (in preparation).  Additionally, we will study the relation between star and cluster formation, and the dependence of environment in Adamo et al. (in preparation)."
    },
    "1402/1402.6186_arXiv.txt": {
        "abstract": "{ The Fluorescence Detector (FD) of the Pierre Auger Observatory provides a nearly calorimetric measurement of the primary particle energy, since the fluorescence light produced is proportional to the energy dissipated by an Extensive Air Shower (EAS) in the atmosphere. The atmosphere therefore acts as a giant calorimeter, whose properties need to be well known during data taking. Aerosols play a key role in this scenario, since their effect on light transmission is highly variable even on a time scale of one hour, and the corresponding correction to EAS energy can range from a few percent to more than 40$\\%$. For this reason, hourly Vertical Aerosol Optical Depth ($\\rm{\\tau_{aer}(h)}$) profiles are provided for each of the four FD stations. Starting from 2004, up to now 9 years of $\\rm{\\tau_{aer}(h)}$ profiles have been produced using data from the Central Laser Facility (CLF) and the eXtreme Laser Facility (XLF) of the Pierre Auger Observatory. The two laser facilities, the techniques developed to measure $\\rm{\\tau_{aer}(h)}$ profiles using laser data and the results will be discussed.} ",
        "introduction": "Ultra High Energy Cosmic Rays (UHECR, $\\rm E>10^{18} eV$) entering the atmosphere cannot be directly detected due to their extremely low flux.  For this reason, the properties of primary particles (energy, mass composition, direction) are deduced from the study of the cascade of secondary particles (Extensive Air Showers, EAS) that originates in the atmosphere due to the interaction of those primaries with air molecules. The Pierre Auger Observatory is the largest detector of EAS ever built, covering an area of $\\rm 3000$ $\\rm km^2$, located in Argentina in the province of Mendoza. The observatory uses two techniques at the same time : the detection of particles at ground level with the Surface Detector (SD) and the observation of the longitudinal development of the EAS by detecting the fluorescence light emitted with the Fluorescence Detector (FD). The SD array is composed of more than 1600 water Cherenkov detectors, overlooked by 27 fluorescence telescopes grouped in 4 sites located at the array periphery. The observatory was completed in 2008. The FD is designed to perform a nearly calorimetric measurement of the energy of cosmic ray primaries: the detected flux of fluorescence photons, emitted by nitrogen air molecules excited by EAS charged particles, is proportional to the energy deposit per unit slant depth of the traversed atmosphere. Due to the constantly changing properties of the calorimeter (i.e. the atmosphere), in which the light is both produced and through which it is transmitted, an extensive system with several instruments has been set up to perform a continuous monitoring of its properties. In particular, the aerosol attenuation of the fluorescence light, highly variable on a time scale of one hour, needs to be constantly measured during data acquisition. If the aerosol attenuation is not taken into account, the shower energy reconstruction is biased by 8 to 25\\% in the energy range measured by the Pierre Auger Observatory. On average, 20\\% of all showers have an energy correction larger than 20\\%, 7\\% of showers are corrected by more than 30\\% and 3\\% of showers are corrected by more than 40\\% \\cite{bib:SegevPaper}. At the Pierre Auger Observatory, hourly vertical aerosol optical depth profiles have been produced for each FD site from January 2004 to December 2012 for a correct reconstruction of FD events. ",
        "conclusions": ""
    },
    "1402/1402.6523_arXiv.txt": {
        "abstract": "{ We propose a supersymmetric scenario in which the small Yukawa couplings for the Dirac neutrino mass term are generated by the spontaneous-breaking of Pecci-Quinn symmetry. In this scenario, a right amount of dark matter relic density can be obtained by either right-handed sneutrino or axino LSP, and a sizable amount of axion dark radiation can be obtained. Interestingly, the decay of right-handed sneutrino NLSP to axino LSP is delayed to around the present epoch, and can leave an observable cosmological background of neutrinos at the energy scale of $\\mathcal{O}(10-100) \\GeV$. }  \\citestyle{plain}  \\begin{document} ",
        "introduction": "The standard model (SM) has been extremely successful in describing subatomic world, but many astrophysical and cosmological observations require a theory beyond SM. One of the apparent shortcomings of SM is the lack of the tiny neutrino mass hinted by atmospheric and solar neutrino oscillations \\cite{Fukuda:1998mi,Ahmad:2001an} (also reactor and long-baseline neutrino oscillations \\cite{Eguchi:2002dm,Michael:2006rx}). The mass-squared differences of neutrino mass eigenstates are now known to be \\cite{Fogli:2012ua} \\beq |\\Delta m_{21}^2| \\simeq 7.5 \\times 10^{-5} \\eV^2, \\quad |\\Delta m_{32}^2| \\simeq 2.4 \\times 10^{-3} \\eV^2 \\eeq This implies that at least one neutrino has a mass of at least $0.05 \\eV$. On the other hand, a recent analysis based on data from Planck satellite mission \\cite{planck} and predictions from other phenomena found a consistent picture of $\\Lambda$CDM model with the sum of the active neutrino masses given by \\cite{Battye:2014} \\beq \\sum m_\\nu = 0.320 \\pm 0.081 \\eV \\eeq When one tries to get such small neutrino masses from the Higgs mechanism with right-handed (RH) neutrinos introduced, very tiny Yukawa couplings of $\\mathcal{O}(10^{-13}-10^{-12})$ are required. % It looks quite puzzling to have such small Yukawa couplings. The best known mechanism for this puzzle is the so-called seesaw mechanism \\cite{GellMann:1980vs,Yanagida:1979as,Mohapatra:1979ia} which uses a large Majorana mass term of RH-neutrinos. However one should note that a Majorana particle has never been observed so far and the small Yukawa couplings may have a dynamical origin.  Besides phenomenological issues, SM also suffers from an esthetic theoretical issue, strong CP problem \\cite{Kim:1986ax,Peccei:2006as} requiring a tuning of $\\mathcal{O}(10^{-10})$ to match experimental data \\cite{Dress:1976bq,Altarev:1996xs}. Axion \\cite{Peccei:1977hh,Peccei:1977ur,Weinberg:1977ma,Wilczek:1977pj,Kim:1979if,Shifman:1979if,Zhitnitsky:1980tq,Dine:1981rt} from the breaking of a global Abelian symmetry (called $U(1)_{\\rm PQ}$) \\cite{Peccei:1977hh,Peccei:1977ur} provides a very simple and attractive solution to this problem, while becoming a good candidate of cold dark matter. Additionally, as we discuss in this paper, the symmetry breaking field associated with $U(1)_{\\rm PQ}$ may be responsible for the small Dirac neutrino mass term.  Meanwhile, low energy supersymmetry (SUSY) is quite attractive because it can provide a fine unification of SM gauge couplings, a natural solution to the hierarchy problem of electroweak Higgs mass and a candidate of dark matter under the assumption of $R$-parity conservation, even though it is facing with \\textit{little hierarchy problem} arising due to the lack of SUSY signature at recent collider experiments \\cite{Feng:2013pwa}. In particular, dark matter might be from the extended non-MSSM sector which is necessary to address various theoretical/phenomenological shortcomings of SM.  In this paper, we propose a supersymmetric extension of SM in which the tiny Yukawa couplings of the Dirac neutrino mass term are generated dynamically by Peccei-Quinn field which breaks $U(1)_{\\rm PQ}$ symmetry spontaneously, and discuss its cosmological implications including dark matter, dark radiation and cosmic neutrino flux.  This paper is organized as follows. In section~\\ref{sec:model}, our model is described. In section~\\ref{sec:cosmo}, cosmological aspects of scalar fields (particularly, RH-sneutrino and saxion fields) are briefly discussed. In section~\\ref{sec:DM}, we estimate abundances of LSP and NLSP before it decays. In section~\\ref{sec:DR}, the dark radiation contribution of relativistic axions produced in the decay of saxion is estimated. In section~\\ref{sec:neu-flux}, a cosmological neutrino flux produced in the decay of NLSP is estimated. In section~\\ref{sec:con}, our conclusion is provided. ",
        "conclusions": "\\label{sec:con} In this paper, we proposed a simple supersymmetric extension of the standard model, in which both the MSSM $\\mu$-term and small Yukawa couplings for the tiny Dirac neutrino mass term are simultaneously generated by the intermediate scale vacuum expectation value of PQ-field which breaks $U(1)_{\\rm PQ}$ symmetry.  It was shown that a right amount of relic density can be obtained by axino LSP produced in the decay of saxion (flaton responsible for thermal inflation) and/or thermally generated neutralinos. The possibility of right-handed sneutrino LSP which can saturate the observed relic density was also discussed. Additionally, it was shown that in the case of thermal inflation a sizable amount of axion dark radiation, which match to the recent data from Planck satellite mission, can also be obtained in a wide range of parameter space.  Interestingly, we found that right-handed sneutrino NLSP decaying to axino LSP can produce a cosmological neutrino flux which may be observable in the near future experiments in the energy range of $\\mathcal{O}(10-100) \\GeV$. This is a very unique possibility of supporting our scenario although there is still a large room of parameter space without any observable signatures."
    },
    "1402/1402.7075_arXiv.txt": {
        "abstract": "\\noindent {\\it Kepler} has identified over 600 multiplanet systems, many of which have several planets with orbital distances smaller than that of Mercury -- quite different from the Solar System.  Because these systems may be difficult to explain in the paradigm of core accretion and disk migration, it has been suggested that they formed in situ within protoplanetary disks with high solid surface densities.  The strong connection between giant planet occurrence and stellar metallicity is thought to be linked to enhanced solid surface densities in disks around metal-rich stars, so the presence of a giant planet can be a detectable sign of planet formation in a high solid surface density disk.  I formulate quantitative predictions for the frequency of long-period giant planets in these in situ models of planet formation by translating the proposed increase in disk mass into an equivalent metallicity enhancement.  I rederive the scaling of giant planet occurrence with metallicity as $P_{\\mathrm{gp}} = 0.05_{-0.02}^{+0.02} \\times 10^{(2.1 \\pm 0.4) [\\mathrm{M/H}]} = 0.08_{-0.03}^{+0.02} \\times 10^{(2.3 \\pm 0.4) [\\mathrm{Fe/H}]}$ and show that there is significant tension between the frequency of giant planets suggested by the minimum mass extrasolar nebula scenario and the observational upper limits.  This fact suggests that high-mass disks alone cannot explain the observed properties of the close-in {\\it Kepler} multiplanet systems and that migration is still a necessary contributor to their formation.  More speculatively, I combine the metallicity scaling of giant planet occurrence with recently published small planet occurrence rates to estimate the number of Solar System analogs in the Galaxy. I find that in the Milky Way there are perhaps $4 \\times 10^{6}$ true Solar System analogs with an FGK star hosting both a terrestrial planet in the habitable zone and a long-period giant planet companion. ",
        "introduction": "{\\it Kepler} has discovered many multiplanet systems with several planets with orbital periods $P < 50$ days.\\footnote{See for example \\citet{bor11a,bor11b}, \\citet{bat13}, and \\citet{bur14}.} Indeed, 40\\% of solar-type stars in the {\\it Kepler} field have at least one planet with $P < 50$ days \\citep[e.g.,][]{fre13}.  Even though these systems differ from the Solar System, their apparent ubiquity suggests that they may represent a frequent outcome of planet formation.  In the traditional minimum mass solar nebula (MMSN) scenario, there is probably insufficient solid material in protoplanetary disks to form the {\\it Kepler} multiplanet systems where they are observed today \\citep{wei77,hay81}.  Instead, formation further out in the parent protoplanetary disk combined with subsequent inward migration has been suggested as one possible formation channel for this class of system \\citep[e.g.,][]{ali06}.  The apparent excess of planets just outside of mean-motion resonances may also support the formation then inward migration scenario \\citep[e.g.,][]{lis11,fab12}.  However, the rate and even direction of migration is known to sensitively depend on the unknown thermodynamic state of the disk \\citep[e.g.,][]{paa10,kle12}.  Alternative models of in situ formation in disks with solid surface densities enhanced beyond the MMSN expectation have also been suggested to explain the ubiquity of close-in multiple systems.  In the minimum mass extrasolar nebula (MMEN) scenario of \\citet{chi13}, the {\\it Kepler} multiplanet systems formed in protoplanetary disks that were about six times more massive than those envisioned in the MMSN scenario.  In that picture, the more massive MMEN disks describe the typical protoplanetary disk in the Galaxy, while the less massive MMSN disk is the outlier. I illustrate the MMEN in Figure~\\ref{fig01}.  On the other hand, \\citet{han12} invoke the rapid inward migration of planetesimals into the inner regions of the disk.  The enhanced solid surface density of the inner disk then naturally leads to the in situ formation of planetary systems closely resembling those observed by {\\it Kepler} \\citep{han13}.  Both models of in situ planet formation described above provide useful, fresh looks at planet formation.  As I will show, both models are also amenable to quantitative tests.  At face value, both models make qualitative predictions for the formation of long-period giant planets. All else being equal, a disk with higher solid surface density in the giant planet forming region has a better chance of forming a giant planet than does a disk with lower solid surface density in the giant planet forming region \\citep[e.g.,][]{lis09}.  Consequently, \\citet{chi13} suggested that the enhanced solid surface density in the MMEN scenario should lead to the efficient formation of giant planets outside of 1 AU. In contrast, the concentration of a significant amount of a disk's solid material in the inner disk as suggested by \\citet{han12} should lead to inefficient formation of giant planets outside of 1 AU.  Unfortunately, it is not currently possible to directly assess with either the transit or radial velocity technique the frequency of long-period giant planets in the observed {\\it Kepler} multiple systems themselves.  However, it is possible to indirectly infer the frequency of long-period giant planets in at least two ways.  First, the frequency can be characterized by proxy, a technique in which objects in the solar neighborhood that can be studied in detail stand in for the more distant {\\it Kepler} objects.  In this case, the close-in multiple systems of low-mass planets discovered in the solar neighborhood with the radial velocity technique are a proxy for the more distant {\\it Kepler} multiple systems.  One can use the published completeness limits of the radial velocity surveys to establish upper limits on the frequency of long-period giant planets in those systems, then compare that upper limit to quantitative predictions of the in situ models of planet formation. Second, it well known that giant planet host stars are preferentially metal enriched \\citep[e.g.,][]{san04,fis05}.  Therefore, the absence of a metallicity effect in solar-type hosts of {\\it Kepler} multiple systems can be used to determine an upper limit on the frequency of giant planets in these systems.  In this paper, I compare statistical upper limits on the frequency of long-period giant planets in the {\\it Kepler} multiple systems with quantitative predictions of the in situ models of planet formation. I find that there is significant tension between the derived upper limits and the expectation from the MMEN scenario, though current samples are not large enough to constrain the \\citet{han12} scenario.  I describe my sample selection in Section 2, I detail my statistical analyses in Section 3, I discuss the results and implications in Section 4, and I summarize my findings in Section 5. ",
        "conclusions": "{\\it Kepler} multiple planet systems with several planets orbiting with periods $P < 50$ days may be difficult to explain in the traditional core accretion and Type I migration paradigm.  In response, two in situ models of plant formation were proposed: the minimum mass extrasolar nebula (MMEN) scenario of \\citet{chi13} and the planetesimal migration scenario of \\citet{han12}.  Both models make predictions for the occurrence rate of long-period giant planets.  In the MMEN scenario, they are ubiquitous.  In the planetesimal migration scenario, they are less common than in the field population.  I find that the prediction from the MMEN scenario for the occurrence of long-period giant planets in multiple low-mass systems discovered with the radial velocity technique fails at 3$\\sigma$.  The lack of metallicity enhancement in the hosts of multiple small-planet systems discovered by {\\it Kepler} provides an independent test of the MMEN scenario, which it also fails at 3$\\sigma$. I am unable to constrain the planetesimal accretion scenario with the current sample.  As a result, migration is still a necessary step in the formation of systems of close-in low-mass planets.  I also rederived the scaling of giant planet occurrence on metallicity, and I find that $P_{\\mathrm{gp}} = 0.05_{-0.02}^{+0.02} \\times 10^{(2.1 \\pm 0.4) [\\mathrm{M/H}]} = 0.08_{-0.03}^{+0.02} \\times 10^{(2.3 \\pm 0.4) [\\mathrm{Fe/H}]}$.  I used these relations to calculate the frequency of ``solar systems\" in the Galaxy, where a solar system is defined as a single FGK star orbited by a planetary system with at least one small planet interior to the orbit of a giant planet.  The presence of a giant planet exterior to an Earth-size planet may be necessary to prevent frequent comet impacts from inhibiting the evolution of life on an otherwise habitable planet.  I find that in the solar neighborhood, about 0.7\\% of solar-type stars have a ``solar system\" consisting of both a Neptune-size planet with $P < 50$ days and a protective long-period companion.  Intriguingly, I find that true Solar System analogs with both a terrestrial planet in the habitable zone and a long-period giant planet companion to protect it occur around only 0.06\\% of solar-type stars. There are perhaps $4 \\times 10^{6}$ such systems in the Galaxy."
    },
    "1402/1402.0180_arXiv.txt": {
        "abstract": "The nonbaryonic dark matter of the Universe can consist of new stable charged species, bound in heavy neutral \"atoms\" by ordinary Coulomb interaction. Stable $\\bar U$ (anti-$U$)quarks of 4th generation, bound in stable colorless ($\\bar U \\bar U \\bar U $) clusters, are captured by the primordial helium, produced in Big Bang  Nucleosynthesis, thus forming neutral \"atoms\" of O-helium (OHe), a specific nuclear interacting dark matter that can provide solution for the puzzles of direct dark matter searches. However, the existence of the 4th generation quarks and leptons should influence the production and decay rates of Higgs boson and is ruled out by the experimental results of the Higgs boson searches at the LHC, if the Higgs boson coupling to 4th generation fermions with is not suppressed. Here we argue that the difference between the three known quark-lepton families and the 4th family can naturally lead to suppression of this coupling, relating the accelerator test for such a composite dark matter scenario to the detailed study of the production and modes of decay of the 125.5 GeV boson, discovered at the LHC. ",
        "introduction": "The cosmological dark matter can consist of dark atoms, in which new stable charged particles are bound by ordinary Coulomb interaction (see \\cite{Levels,Levels1,mpla,DMRev,4QGC} for review and references). In order to avoid anomalous isotopes overproduction, stable particles with charge -1 (and corresponding antiparticles as tera-particles \\cite{Glashow}) should be absent \\cite{Fargion:2005xz}, so that stable negatively charged particles should have charge -2 only.  There exist several types of particle models, in which heavy stable -2  charged species, $O^{--}$, are predicted: \\begin{itemize} \\item[(a)] AC-leptons, predicted in the extension of standard model, based on the approach of almost-commutative geometry \\cite{Khlopov:2006dk,5,FKS,bookAC}. \\item[(b)] Technileptons and anti-technibaryons in the framework of walking technicolor models (WTC) \\cite{KK,Sannino:2004qp,Hong:2004td,Dietrich:2005jn,Dietrich:2005wk,Gudnason:2006ug,Gudnason:2006yj}. \\item[(c)] and, finally, stable \"heavy quark clusters\" $\\bar U \\bar U \\bar U$ formed by anti-$U$ quark of 4th generation \\cite{Khlopov:2006dk,Q,I,lom,KPS06,Belotsky:2008se}. \\end{itemize}  All these models also predict corresponding +2 charge particles. If these positively charged particles remain free in the early Universe, they can recombine with ordinary electrons in anomalous helium, which is strongly constrained in the terrestrial matter. Therefore cosmological scenario should provide a  mechanism, which suppresses anomalous helium. There are two possibilities, requiring two different mechanisms of such suppression: \\begin{itemize} \\item[(i)] The abundance of anomalous helium in the Galaxy may be significant, but in the terrestrial matter there exists a recombination mechanism suppressing this abundance below experimental upper limits \\cite{Khlopov:2006dk,FKS}. The existence of a new strict U(1) gauge symmetry, causing new Coulomb like long range interaction between charged dark matter particles, is crucial for this mechanism. Therefore the existence of dark radiation in the form of hidden photons is inevitable in this approach. \\item[(ii)] Free positively charged particles are already suppressed in the early Universe and the abundance of anomalous helium in the Galaxy is negligible \\cite{mpla,DMRev,I}. \\end{itemize} These two possibilities correspond to two different cosmological scenarios of dark atoms. The first one is realized in the scenario with AC leptons, forming neutral AC atoms \\cite{FKS}. The second assumes charge asymmetric case with the excess of $O^{--}$, which form atom-like states with primordial helium \\cite{mpla,DMRev,I}.   If new stable species belong to non-trivial representations of electroweak SU(2) group, sphaleron transitions at high temperatures can provide the relationship between baryon asymmetry and excess of -2 charge stable species, as it was demonstrated in the case of WTC \\cite{KK,KK2,unesco,iwara}.  After it is formed in the Standard Big Bang Nucleosynthesis (SBBN), $^4He$ screens the $O^{--}$ charged particles in composite $(^4He^{++}O^{--})$ {\\it O-helium} ``atoms'' \\cite{I}.  In all the models of O-helium, $O^{--}$ behaves either as lepton or as specific \"heavy quark cluster\" with strongly suppressed hadronic interaction. Therefore O-helium interaction with matter is determined by nuclear interaction of $He$. These neutral primordial nuclear interacting objects can explain the modern dark matter density and represent a nontrivial form of strongly interacting dark matter \\cite{McGuire:2001qj,McGuire1,McGuire2,Starkman,Starkman2,Starkman3,Starkman4,Starkman5,Starkman6}.  The cosmological scenario of O-helium Universe allows to explain many results of experimental searches for dark matter \\cite{mpla}. Such scenario is insensitive to the properties of $O^{--}$, since the main features of OHe dark atoms are determined by their nuclear interacting helium shell. It challenges experimental probes for the new stable charged particles at accelerators and such probes strongly depend on the nature of $O^{--}$.  Stable $-2$ charge states ($O^{--}$) can be elementary like AC-leptons or technileptons, or look like elementary as technibaryons. The latter, composed of techniquarks, reveal their structure at much higher energy scale and should be produced at the LHC as elementary species. They can also be composite like \"heavy quark clusters\" $\\bar U \\bar U \\bar U$ formed by anti-$U$ quark in the models of stable fourth generation \\cite{Q,I}.   In the context of composite dark matter scenario accelerator probe for new stable quark-lepton generation acquires the meaning of critical test for the existence of charged constituents of cosmological dark matter.  The signature for double charged AC leptons and techniparticles is unique and distinctive what has already allowed to obtain the lower bound on their mass of 430 GeV in the ATLAS experiment \\cite{2qatlas}.  Test for composite $O^{--}$ at the LHC can be only indirect: through the search for heavy hadrons, composed of single $U$ or $\\bar U$ and light quarks (similar to R-hadrons) \\cite{4QGC}, or by virtual effects of 4th generation fermions in the processes with known particles. Here we study a possibility for experimental probe of the hypothesis of stable 4th generation in the studies of 125.5 GeV Higgs boson, discovered in the ATLAS \\cite{atlas,atlas1} and the CMS experiments \\cite{cms1,cms} at the LHC. The results of these studies \\cite{atlas,atals1,cms1,cms} indicate that the number of the detected events, being the production cross section times the decay rate of Higgs boson to two-photon channel, is consistent with the prediction of the Standard model. On the other hand, as it was first revealed in \\cite{nuHiggs} the existence of 4th generation leads to enhancement of the main mechanism of Higgs boson production in $pp$ collisions, what puts constraints on the effect of 4th generation particles and practically excludes the possibility of their full strength coupling to 125.5 GeV Higgs boson.  In the model of stable 4th generation the difference of these fermions from the quarks and leptons of the three known families is related to some new conserved charge (which can be even a gauge charge) that protects the stability the lightest quarks and leptons ($U$-quark and the 4th neutrino $N$). The experimental lower limits on the new quarks and leptons make these particles to be heavier than the three light families, what can be explained by the existence of an additional mechanism of their mass generation, e.g. in the framework of multi-Higgs models. It can naturally lead to suppression of the coupling of 4th generation fermions to the 125.5 GeV Higgs boson, discovered at the LHC. Here we explore a possibility to make the 4th generation hypothesis consistent with the experimental data on the two gamma decays of Higgs boson, what opens the door to the indirect probes of the charged constituents of composite dark matter in the detailed studies of production and modes of decay of the 125.5 GeV Higgs boson. ",
        "conclusions": "The cosmological dark matter can be formed by stable heavy double charged particles bound in neutral OHe dark atoms with primordial He nuclei by ordinary Coulomb interaction. This scenario sheds new light on the nature of dark matter and offers nontrivial solution for the puzzles of direct dark matter searches. It can be realized in the model of stable 4th generation and challenges for its experimental probe at accelerators. In the context of this scenario search for effects of new heavy quarks and leptons acquires the meaning of direct experimental probe for charged constituents of dark atoms of dark matter.  The $O^{--}$ constituents of OHe in the model of stable 4th generation are \"heavy antiquark clusters\" $\\bar U \\bar U \\bar U$. Production of such clusters (and their antiparticles) at accelerators is virtually impossible. Therefore experimental test of the hypothesis of stable 4th generation is reduced to the search for stable or metastable $U$ hadrons, containing only single heavy quark or antiquark, or to the studies of virtual effects of 4th generation quarks and leptons in the processes with known particles.  The discovery of the 125.5 GeV Higgs boson at the LHC opens the new room for such indirect test of the model of stable 4th generation. The number of detected events of decays of this boson to the two-photon channel is consistent within the experimental errors with the prediction of the Standard model, putting severe constraints on the contribution of new quarks and leptons. On the other hand, the existence of heavy quarks of the 4th generation should lead to enhancement of the gluon fusion mechanism of Higgs boson production, which is its dominant production mechanism in $pp$ collisions. Therefore, to be compatible with the experimental data the model of the stable 4th generation should involve a mechanism of suppression of new quark and lepton couplings to the 125.5 GeV Higgs boson. Taking into account the difference of the 4th generation from the three known families of quarks and leptons and in particular the lower limits on the masses of new quarks and leptons, one can assume some additional mechanisms of their mass generation, what can lead to suppression of their couplings to the 125.5 GeV Higgs boson.  In the present work we have shown that, indeed, the suppression in the Higgs boson couplings to 4th generation quarks and lepton the can make compatible the existence of this generation with the experimental data on the two-photon decays of the 125.5 GeV Higgs boson. We consider this result as the first step in the thorough investigation of the predictions of the model of stable 4th generation for the whole set of decay channels of Higgs boson. The confrontation of these predictions with the detailed experimental study of the 125.5 GeV Higgs boson will provide the complete test for the composite dark matter scenarios based on the model of the stable 4th generation."
    },
    "1402/1402.2466_arXiv.txt": {
        "abstract": "We study the time evolution of the impact probability for synthetic but realistic impacting and close approaching asteroids detected in a simulated all-sky survey.  We use the impact probability to calculate the impact warning time ($t_{warn}$) as the time interval between when an object reaches a Palermo Scale value of $-2$ and when it impacts Earth.  A simple argument shows that $t_{warn} \\propto D^x$ with the exponent in the range $[1.0,1.5]$ and our derived value was $x=1.3 \\pm 0.1$ .  The low-precision astrometry from the single simulated all-sky survey could require many days or weeks to establish an imminent impact for asteroids $>100\\meter$ diameter that are discovered far from Earth.  Most close approaching asteroids are quickly identified as not being impactors but a size-dependent percentage, even for those $>50\\meter$ diameter, have a persistent impact probability of $>10^{-6}$ on the day of closest approach. Thus, a single all-sky survey can be of tremendous value in identifying Earth impacting and close approaching asteroids in advance of their closest approach but it can not solve the problem on its own: high-precision astrometry from other optical or radar systems is necessary to rapidly establish an object as an impactor or close approacher.  We show that the parallax afforded by surveying the sky from two sites is only of benefit for a small fraction of the smallest objects detected within a couple days before impact: probably not enough to justify the increased operating costs of a 2-site survey. Finally, the survey cadence within a fixed time span is relatively unimportant to the impact probability calculation.  We tested three different reasonable cadences and found that one provided $\\sim10\\times$ higher (better) value for the impact probability on the discovery night for the smallest ($10\\meter$ diameter) objects but the consequences on the overall impact probability calculation is negligible. ",
        "introduction": "The past quarter century has witnessed an exponential increase in the number of known near-Earth objects\\footnote{NEOs are asteroids or comets that have a perihelion distance $<1.3\\au$} (NEO) accompanied by a concomitant improvement in the ability to calculate their impact probabilities with Earth.  The job of identifying the largest and most hazardous NEOs, those larger than about $1\\km$ diameter, has mostly been accomplished and asteroid surveys are now focussing on the individually less hazardous but far more numerous smaller asteroids.  The large asteroids can be detected at great distances years to centuries in advance of their impact but the smaller asteroids may only be detected on their final approach, if at all, since about 40\\% of them must approach from the direction of the Sun in daylight sky.  This work quantifies how the impact probability and warning time evolve in the impact apparition for the smaller asteroids as a function of their size, time after discovery, and observing cadence.  In particular, we examine whether the parallax afforded by observations at nearly the same time from two independent observatories provides leverage in improving the impact probability calculation or increasing the impact warning time.  The Catalina Sky Survey \\citep[CSS,\\ ][]{Larson1998} and \\PSone\\ system \\citep[Pan-STARRS; \\eg][]{Kaiser2002,Hodapp2004} currently dominate the field of NEO discovery --- almost 90\\% of all NEOs and about 75\\% of all potentially hazardous objects\\footnote{PHOs are NEOs that have a minimum orbital intersection distance \\citep[\\eg][]{Gronchi2005} with Earth of $<0.05\\au$ and absolute magnitude $H<22$} (PHO) were discovered by these two surveys in calendar years 2012 and 2013.  The known population of NEOs larger than $1\\km$ diameter is $>90$\\% complete \\citep{Mainzer2011c} so the discovery rate of NEOs in this size range has decreased by about a factor of 6 from a peak of 93 in the year 2000 to about 15/year in the last two years.  Despite the success of the surveys in the past few decades it remains the case that the most likely warning time for an impact is {\\it zero} --- contemporary surveys are unlikely to detect smaller but still dangerous asteroids because they do not survey the entire sky deeply or regularly enough to identify the next impactor.  The surveys are further limited by the simple fact that ground-based facilities can not survey during the day and about 40\\% of all impactors will approach from the direction of the Sun.  These problems were spectacularly highlighted by the Chelyabinsk impact on the morning of 15 February 2013 \\citep[\\eg][]{Brown2013,Borovicka2013} --- with absolutely no warning a $\\sim17\\meter$ diameter object blew up in the atmosphere with an energy equivalent to about 500\\,kilotons of TNT, damaging buildings 50\\,km away in the city of Chelyabinsk and injuring about 1,500 people.  The impact risk associated with the unknown objects $>1\\km$ diameter is now comparable to the impact risk with the much more numerous but individually less destructive objects with diameters $<1\\km$.  The new balance in the impact risk, along with the realization that smaller impacts may be more numerous but less destructive than anticipated a decade ago \\citep[\\eg][]{Brown2013}, has contributed to an increased interest and funding for the NEO survey programs in recent years. \\eg\\ NASA's NEO Observations (NEOO) program office now solicits\\footnote{ROSES 2011 NEOO solicitation section C.9.1.1.} proposals for surveys that `provide capability to detect the subset of 90\\% of PHOs down to 140 meters in size'  The smaller NEOs are more difficult to detect than the larger ones, will be detected ({\\it if} they are detected) closer to Earth, and consequently have shorter observational arcs.  The limited time range of the set of detections can make it difficult to identify real impactors even during the apparition in which the impact will take place.  This was not the case for the few-meter diameter asteroids \\TC \\citep[\\eg][]{Jenniskens2009} and 2014~AA\\footnote{Minor Planet Electronic Circular 2014-A02}, the only natural objects to be discovered before striking Earth.   The very smallest objects will be discovered so close to Earth that, if individual detections of the object can be associated with one another as a `tracklet' \\citep{Denneau2013}, the non-linear motion of the detections on the sky-plane due to topocentric parallax can provide enough leverage in the orbit solution to predict an impact.  The observable characteristics of NEOs that will impact Earth can be quite different from those of other NEOs \\citep[\\eg][]{Farnocchia2012,Veres2009,Chesley2004}.  For instance, their observable steady-state distribution on the sky-plane is a function of their size and time before impact.  Decades before impact they tend to be concentrated in `sweet spots' near the ecliptic and within about $120\\arcdeg$ of the Sun.  As the time until impact decreases from weeks to days they spread out over most of the sky but there are still concentrations in the direction looking towards and away from the Sun.  An object on its `death plunge' must be moving directly towards Earth in a geocentric reference frame so that about a week before impact its apparent rate of motion may be small --- likely mimicking the rate of motion of much more distant and totally harmless asteroids and perhaps not triggering followup that would allow an impact probability calculation.  The techniques employed for the impact probability calculation have evolved dramatically over the past few decades with the realization that asteroid impacts have shaped the Moon's surface and influenced the evolution of life on Earth.  Indeed, it was only 34 years ago that \\citet{Alvarez1980} proposed that the KT extinction was the result of an asteroid impact and, even though \\citet{Opik1952} stated that ``Over a dozen meteor craters are at present known on the earth's surface'', it was only in 1960 that \\citet{Chao1960} found strong physical evidence that Meteor Crater in Arizona, USA, was formed in an impact event.  \\citet{Opik1952}'s estimated collision rates using the `Theory of Probabilities' for the entire NEO population were surprisingly good given that only six NEOs were known at the time.  His collision probability formulae formed the basis of much of the impact collision work in the next decades \\citep[\\eg][]{Bottke1996,Kessler1981} but were eventually supplanted by new numerical techniques \\citep{Milani2002}.  The two primary operational asteroid impact warning systems, the Jet Propulsion Laboratory's Sentry system and the NEODyS CLOMON2 system, calculate the collision probability by generating synthetic `Virtual Asteroids' (VA) on orbits that are consistent with the known set of observations and propagating all of them into the future with a N-body integrator to search for impacts \\citep{Milani2005}.  These impact warning systems are based on a geometric sampling technique for which the identification of the Virtual Impactors (VI) is performed on the line-of-variation \\citep[LOV,\\ ][]{Milani2005b} thus avoiding the poor efficiency inherent to the Monte Carlo methods, especially when the collision probability is small.  Impact predictions are extremely sensistive to the orbit accuracy which depends on many factors but the primary drivers are the length of the observational arc and the astrometric accuracy \\citep[\\eg\\ ][]{Desmars2013}.  The longer the arc and the better the astrometry the more accurate the orbit.  The latter effect is best illustrated by radar detection of asteroids that provide exquisite range and range-rate information thereby dramatically improving the impact probability accuracy and/or extending the time frame during which the impact probability can be calculated \\citep{Ostro2002}.  Impactors can be either direct or resonant \\citep{Milani1999}.  Direct impactors collide with Earth during their first known encounter and must be discovered far away to have a large warning time.  The warning time for small impactors can be significantly less than one orbital period because they have to be close to Earth to be detected.  On the other hand, resonant impactors experience intervening Earth encounters before collision.  The intervening encounters are the main source of non-linearity in the dynamics and usually prevent a conclusive assessment of the impact threat but provide additional observational opportunities to detect and constrain their orbits and the impact threat.  In this work we focus on the evolution of the collision probability with time for a single survey and concentrate on direct Earth impactors that are detected in the apparition during which the impact occurs.  The smaller the asteroid the more likely this scenario as the likelihood that small asteroids will be detected in earlier apparitions is 1) small and, even if they are detected, 2) it is unlikely that they will be recoverable in future apparitions because of the large uncertainties in their ephemeris based on the short observational arcs in the discovery apparition.  Thus, we concentrate on collision probability evolution with time for $300\\meter$, $100\\meter$, $50\\meter$ and $10\\meter$ diameter impactors.  We also explore whether the collision probability calculation benefits from simultaneous or nearly-simultaneous parallax measurements from two observatories.  The heliocentric motion of the impactor and Earth as well as the topocentric rotation of the observer about the geocenter produce `parallax' between successive observations of the same object even from the same site. For very close objects that will impact within days of discovery there may be benefits from the two-site scenario --- especially in rapidly identifying the object as an impactor.  Finally, we measure the single-system impact warning time as a function of impactor diameter.  We expect the warning time to be longer for larger objects but the exact relationship between diameter and warning time is not intuitively obvious.  The larger objects are discovered at greater distances where their rate of motion is similar to the much more distant main belt objects and the impact probability will be much smaller.  If the impact probability is too small it may not cross the threshold to flag the object as an imminent impactor. ",
        "conclusions": "We have performed a simulation of a single all-sky asteroid survey to study the time evolution of the calculated impact probability for both real and false impactors.  We also studied the utility of using two observatories at different locations to perform the survey to take advantage of the parallactic displacement in the detections of the same object.  As expected, the impact probability for impactors typically increases monotonically with time after discovery and is larger at the time of discovery for small objects that are detected closer to Earth and with less time to impact.  We found that the impact warning time, the time interval between when the impact probability reaches -2 on the Palermo Scale and when the impact takes place, increases with diameter according to $t_{warn} \\propto D^{1.3}$ and developed a simple mathematical argument that the exponent should be in the range $[1.0,1.5]$.  Close approaching asteroids can almost always be unambiguously identified as non-impactors but a small percentage will have a non-zero impact probability even on the day of (false) impact. \\ie\\ the simulated survey on its own is unable to eliminate the impact risk.  The fraction of objects for which a persistent impact risk exists at the time of impact increases with decreasing diameter of the object because the small objects have smaller observational arc lengths and concomitantly less precise orbit elements.  The combination of the PHO size-frequency distribution with the probability of detecting false impactors suggests that the single all-sky system alone will generate 100s of potential impactors that must be ruled out with other followup facilities.  The calculated impact probability can take surprisingly long to reach $\\sim100$\\% with just the results from a single low-precision astrometric survey.  The impact probability may reach 100\\% only a few days before impact even for $300\\meter$ diameter objects detected a month in advance and imaged nightly thereafter .  Our simulations suggest that a 2-site survey is unnecessary, at least in terms of the incremental benefit in improving the impact probability calculation.  The parallax afforded by this scenario only improves the impact probability calculation for a small fraction of the smallest asteroids detected shortly before impact.  The 2-site survey offers many different cadence options and some can provide more efficient impact probability calculations than others.  The derived impact probability was $\\sim10\\times$ higher (\\ie\\ better) on the discovery night using the `full-shift' cadence compared to the other two cadences.  This suggests that a real survey that implements the 2-site scenario should carefully test different cadences to select one that maximizes the efficiency and accuracy of the impact probability on the discovery night.  The effect of survey cadence on the impact probability calculation is negligible on successive nights."
    },
    "1402/1402.6842_arXiv.txt": {
        "abstract": "We present a sample of normal type~Ia supernovae from the Nearby Supernova Factory dataset with spectrophotometry at sufficiently late phases to estimate the ejected mass using the bolometric light curve. We measure \\nickel\\ masses from the peak bolometric luminosity, then compare the luminosity in the \\cobalt-decay tail to the expected rate of radioactive energy release from ejecta of a given mass.  We infer the ejected mass in a Bayesian context using a semi-analytic model of the ejecta, incorporating constraints from contemporary numerical models as priors on the density structure and distribution of \\nickel\\ throughout the ejecta. We find a strong correlation between ejected mass and light curve decline rate, and consequently \\nickel\\ mass, with ejected masses in our data ranging from 0.9--1.4~\\Msol. Most fast-declining \\mbox{(SALT2 $x_1 < -1$)} normal SNe~Ia have significantly sub-Chandrasekhar ejected masses in our fiducial analysis. ",
        "introduction": "Type Ia supernovae (SNe~Ia) have been used for well over a decade as precision luminosity distance indicators, leading to the discovery of the universe's accelerated expansion \\citep{riess98,scp99} \\revised{which has been measured in contemporary studies with increasing precision} \\citep{hicken09,kessler09,sullivan11b,suzuki12}. SN~Ia luminosities can be measured to an accuracy of $\\sim 0.15$~mag using correlations between the luminosity, colour, and light curve width \\citep{riess96,tripp98,phillips99,goldhaber01}, and many recent and ongoing studies have sought to further reduce this dispersion by looking for new correlations between SN~Ia luminosities and their spectroscopic properties \\citep{sjb09,wang09,csp10,fk10}.  The spectra of SNe~Ia show no hydrogen, no helium, and strong intermediate-mass element signatures; they are generally understood to be thermonuclear explosions of carbon/oxygen white dwarfs in binary systems.  The absence of a detectable shock breakout in the early light curve of the nearby SN~Ia~2011fe \\citep{nugent11,bloom12} provides direct evidence that the progenitor primary must be a compact object such as a white dwarf.  However, many variables remain which can affect the explosion, including the evolutionary state of the white dwarf progenitor's binary companion, the circumstellar environment, the explosion trigger, and the progress of nuclear burning in the explosion. The low luminosities, small radii, and relatively clean environments of white dwarfs make SN~Ia progenitor systems notoriously hard to constrain. Uncovering the nature of SN~Ia progenitor systems and explosions is therefore an interesting puzzle in its own right.  From a cosmological viewpoint, if two or more SN~Ia progenitor channels exist which have slightly different peak luminosities or luminosity standardization relations, and their relative rates evolve with redshift, the resulting shift in the mean luminosity could mimic a time-varying dark energy equation of state \\citep{linder06}.  The two main competing SN~Ia progenitor scenarios are the \\emph{single-degenerate} scenario \\citep{wi73}, in which a carbon/oxygen white dwarf slowly accretes mass from a non-degenerate companion until exploding near the Chandrasekhar mass, and the \\emph{double-degenerate} scenario \\citep{it84}, in which two white dwarfs collide or merge.  The classical formulations of these scenarios assume the primary white dwarf must explode near the Chandrasekhar limit; however, in the \\emph{sub-Chandrasekhar double-detonation} variant, a sub-Chandrasekhar-mass white dwarf can be made to explode by the detonation of a layer of helium on its surface, accreted from the binary companion \\citep{ww94,sim10,fink10,kromer10,sim12}. Distinguishing which of these models accounts for the majority of spectroscopically ``normal'' \\citep{bfn93}, hence cosmologically useful, SNe~Ia has been a very active subject of current research \\citep[for a recent review see][]{wh12}. Binary population synthesis models of the Chandrasekhar-mass single-degenerate and double-degenerate channels often have trouble producing enough SNe~Ia to reproduce the observed rate \\revised{\\citep[but see][]{hp04,ruiter11}}; this is one of the main motivations for investigating sub-Chandrasekhar models \\citep{vkcj10}.  The mass of the progenitor is a fundamental physical variable with power to differentiate between different progenitor scenarios. While Chandrasekhar-mass delayed detonations have been historically favored, viable super-Chandrasekhar-mass evolution pathways and explosion models have been proposed for both single-degenerate \\citep{justham11,hachisu11,rds12} and double-degenerate \\citep{pakmor10,pakmor11,pakmor12} SN~Ia progenitors, and sub-Chandrasekhar-mass models must necessarily involve a different explosion trigger than any of these.  The white dwarf progenitor is totally disrupted in theoretical models of normal SNe~Ia, although a bound remnant may remain in some models which try to reproduce underluminous, peculiar events such as SN~2002cx \\citep{kromer13}.  For normal SNe~Ia, then, measuring the progenitor mass reduces to measuring the ejected mass. Nebular-phase spectra can be used to estimate the mass of iron-peak elements in the ejecta \\citep[e.g.][]{mazzali07}, but only the closest SNe~Ia are bright enough to yield high-quality spectra in nebular phase $\\sim 1$~year after explosion, which limits the number of SNe on which this technique can be used.  \\citet{stritz06} used SN~Ia quasi-bolometric light curves ($UBVRI$) in early nebular phase (50--100~days after $B$-band maximum light) to estimate the ejected mass, as follows:  The mass of \\nickel, \\revised{the radioactive decay of which} powers the near-maximum light curve of normal SNe~Ia, can be inferred from the bolometric luminosity at maximum light \\citep{arnett82}.  The decay of \\cobalt, itself a decay product of \\nickel, powers the post-maximum light curve. At sufficiently late times, the shape of the bolometric light curve is sensitive to the degree of trapping of gamma rays from \\cobalt\\ decay \\citep{jeffery99}; greater ejected masses provide greater optical depth to Compton scattering, and hence higher luminosity, for a given phase and \\nickel\\ mass. \\citet{scalzo10,scalzo12} refined this method by including more accurate near-infrared (NIR) corrections and a set of prior constraints on model inputs from contemporary explosion models, using it to estimate the masses of several candidate super-Chandrasekhar-mass SNe~Ia; they found ejected masses of $2.30^{+0.27}_{-0.24}$~\\Msol\\ for the superluminous SN~Ia~2007if and $1.79^{+0.28}_{-0.21}$~\\Msol\\ for the spectroscopically 1991T-like \\snf{080723-012}, interpreting them as double-degenerate explosions powered entirely by radioactive decay.  In the current work, we use this method as implemented in \\citet{scalzo12} on a set of \\emph{normal} SNe~Ia, attempting to quantify the distribution of progenitor mass scales in the context of different progenitor scenarios. Our supernova discoveries, our sample selection, and the provenance of our data are described in \\S\\ref{sec:observations}. Our \\revised{method for constructing} full $UBVRIYJHK$ (3300--\\revised{23900}~\\AA) bolometric light curves for 19 spectroscopically normal SNe~Ia, (including NIR corrections for the $YJHK$ flux which we do not observe), are presented in \\S\\ref{sec:analysis}.  We briefly review the assumptions of our ejected mass reconstruction method in \\S\\ref{sec:modeling}, and present the reconstructed masses for our 19 SNe.  We also present \\revised{ejected mass and \\nickel\\ mass reconstructions based on} synthetic observables from a series of contemporary explosion models. In \\S\\ref{sec:discussion} we examine correlations between \\revised{ejected mass} and other quantities, such as photospheric light curve fit parameters (decline rate and colour) and \\nickel\\ \\revised{mass}. We summarize and conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": ""
    },
    "1402/1402.3460_arXiv.txt": {
        "abstract": "The IceCube Neutrino Observatory is a large Cherenkov detector instrumenting $1\\,\\mathrm{km}^3$ of Antarctic ice. The detector can be used to search for signatures of particle physics beyond the Standard Model. Here, we describe the search for non-relativistic, magnetic monopoles as remnants of the GUT (\\textbf{G}rand \\textbf{U}nified \\textbf{T}heory) era shortly after the Big Bang. Depending on the underlying gauge group these monopoles may catalyze the decay of nucleons via the Rubakov-Callan effect with a cross section suggested to be in the range of $10^{-27}\\,\\mathrm{cm^2}$ to $10^{-21}\\,\\mathrm{cm^2}$. In IceCube, the Cherenkov light from nucleon decays along the monopole trajectory would produce a characteristic hit pattern. This paper presents the results of an analysis of first data taken from May 2011 until May 2012 with a dedicated slow-particle trigger for DeepCore, a subdetector of IceCube. A second analysis provides better sensitivity for the brightest non-relativistic monopoles using data taken from May 2009 until May 2010. In both analyses no monopole signal was observed. For catalysis cross sections of $10^{-22}\\,(10^{-24})\\,\\mathrm{cm^2}$ the flux of non-relativistic GUT monopoles is constrained up to a level of $\\Phi_{90} \\le 10^{-18}\\,(10^{-17})\\,\\mathrm{cm^{-2}s^{-1}sr^{-1}}$ at a 90\\% confidence level, which is three orders of magnitude below the Parker bound. The limits assume a dominant decay of the proton into a positron and a neutral pion. These results improve the current best experimental limits by one to two orders of magnitude, for a wide range of assumed speeds and catalysis cross sections. ",
        "introduction": "\\label{sec:intro}  Magnetic monopoles are particles carrying a quantized magnetic charge and are predicted in various theories. In classical electrodynamics, their existence would symmetrize Maxwell's equations with respect to the sources of the electromagnetic field. Quantum mechanically, the existence of magnetic monopoles implies that both electric charge and the hypothetical magnetic charge, are quantized, given that the associated electromagnetic fields still satisfy Maxwell's equations \\cite{Dirac:1931}. The resulting magnetic elementary charge, called the Dirac charge $g_{\\mathrm{D}}$, is  \\begin{equation} g_{\\mathrm{D}}=\\frac{e}{2\\alpha}, \\end{equation} where $e$ is the electric elementary charge and $\\alpha$ is the fine structure constant.  In Grand Unified Theories (GUTs) \\cite{Georgi:1974FirstGUT} magnetic monopoles appear as stable, finite energy solutions of the field equations \\cite{tHooft:1974MonopoleSolution,Polyakov:1974MonopoleSolution}. The predicted masses range from $10^5\\,\\mathrm{GeV}$ to $10^{17}\\,\\mathrm{GeV}$ \\cite{Georgi:1974_Interaction_GUT,Daniel:1979_SU5Monopoles,Lazarides,Kephart,Wick:CosmicFluxMM} and their magnetic charges are integer multiples of the Dirac charge  $\tg_{\\mathrm{D}} $. The lower part of the mass range up to $\\sim 10^{13}\\,\\mathrm{GeV}$ refers to intermediate mass monopoles (IMMs) which arise from intermediate stages of symmetry breaking below the GUT scale. In contrast the superheavy monopoles with masses at the GUT scale may have been created during the phase transition associated with the spontaneous breakdown of the unified gauge symmetry in the early universe at $\\sim 10^{-36}\\,\\mathrm{s}$ after the Big Bang \\cite{Kibble}. The monopole mass and charge depend on the underlying gauge group, the symmetry breaking hierarchy, and the type and temperature of the phase transition in a particular GUT.  Since magnetic monopoles are stable, they should still be present in cosmic rays. The number density today depends on the existence of an inflationary epoch and on the time of creation, which could be before, during or after this epoch \\cite{Preskill:MM}. Since then, monopoles have been accelerated by large-scale cosmic magnetic fields. The kinetic-energy gain by passing through a magnetic field $B$ is given by  \\begin{equation} E_{\\mathrm{kin}}=g\\int\\limits_{\\mathrm{path}} \\vec{B} \\cdot d\\vec{l}, \\end{equation} where $g =n\\cdot g_{\\mathrm{D}} $ is the magnetic charge. The maximum kinetic energy of a magnetic monopole due to acceleration in cosmic magnetic fields is rather uncertain but can reach $\\sim 10^{14}\\,\\mathrm{GeV}$ \\cite{Wick:CosmicFluxMM}. Therefore, monopoles with masses at, or above, this energy scale should be non-relativistic. Based on the propagation of magnetic monopoles in the Galactic magnetic field an upper bound on the monopole flux can be calculated, assuming the Galactic magnetic field does not decrease faster than it can be regenerated. This assumption constrains the monopole flux to be less than $10^{-15}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}\\,\\mathrm{sr}^{-1}$, which is called the Parker Bound \\cite{Turner:ParkerBound,Groom:SearchSupermassiveMM}. Taking into account the fields during galaxy formation, the limit was extended by Adams et al. to be less than $10^{-16}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}\\,\\mathrm{sr}^{-1}$ for monopoles with masses below $10^{17}\\,\\mathrm{GeV}$ \\cite{ExtensionParkerBound}.  Many experiments have searched for relic magnetic monopoles, but there is no experimental proof for their existence. The current best limits for magnetic monopoles constrain their flux to a level of $\\sim 10^{-16}-10^{-18}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}\\,\\mathrm{sr}^{-1}$ depending on the monopole speed and interaction mechanism \\cite{FinalMACRO,MACRO:SlowMonopoles,icecube:relMPs,IceCube:ICRC2013_relMPs}. Consequently, searches for magnetic monopoles require very large detectors.  The IceCube Neutrino Observatory currently is the world's largest neutrino detector. The primary goal is the detection of Cherenkov light from electrically-charged secondary particles produced when high-energy astrophysical neutrinos interact in the surrounding matter \\cite{IceCubeSensitivity}. However, IceCube can also be used to search for magnetic monopoles. Depending on their speed monopoles have different signatures in  IceCube. Relativistic monopoles with a speed above the Cherenkov threshold, e.g. $\\beta \\approx 0.76$ in ice, can be detected by the Cherenkov light they directly produce \\cite{Tompkins:CherenkovEmission}. Non-relativistic monopoles that catalyze the decay of nucleons in the detector medium can, in contrast, be detected by the Cherenkov light from electrically charged secondary particles produced in subsequent nucleon decays along the monopole trajectory (Sec. \\ref{subsec:Rubakov-Callan_Effect}). Therefore, different analysis strategies are needed in order to cover both detection channels. This paper presents the results of a search for non-relativistic magnetic monopoles which would catalyze the proton decays via the Rubakov-Callan effect in IceCube. ",
        "conclusions": "Data taken from May 2011 until May 2012 with a dedicated slow-particle trigger and for the brightest monopoles data taken from May 2009 until May 2010 with standard-IceCube-triggers were analysed. The analysis, which is based on data of the slow-particle trigger, was developed by using simulated monopole events and experimental data to estimate background properties. For this first analysis of such a signal in IceCube a robust approach based on a single final selection criterion and the comparison between the number of expected background events and observed experimental events is chosen. Using experimental data, the number of expected background events can be estimated to $n_{\\mathrm{b}} = 3.2^{+1.8}_{-1.1}$.  The IC-59 analysis based on standard-IceCube-triggers is sensitive only for bright monopoles with $\\sigma_{\\mathrm{cat}} > 1.7 \\cdot 10^{-23}\\,\\mathrm{cm^2}$. The analysis used Boosted Decision Trees (BDT) to discriminate between monopole signal and background. The expected number of background events is derived from a fit of the BDT scores tails with an exponential function for each ($\\beta$,$\\lambda_{\\mathrm{cat}}$). The number of observed events after unblinding is $1$ for an expected background of $4.8_{-0.6}^{+0.7}$. This event contains multiple coincident muons, which renders it compatible with a background event. The obtained flux limits for $\\beta=10^{-2}$ and $\\lambda_{\\mathrm{cat}} =0.01$\\,m, $0.001$\\,m from the IC-59 analysis are better than the ones from the IC-86/DeepCore analysis because of the bigger effective area. For $\\beta=10^{-3}$ the limits are comparable since the standard IceCube triggers are less sensitive to the monopole signal in comparison to the dedicated slow-particle trigger.  In both analyses no monopole signal has been observed. Thus, the limits on the flux of non-relativistic magnetic monopoles -- catalysing the proton decay -- are improved by about more than one order of magnitude in comparison to MACRO \\cite{MACRO:SlowMonopoles} for most of the investigated parameter space and reach down to about three orders of magnitude below the Parker limit.  Since May 2012 the dedicated slow-particle trigger has been updated to the full IceCube detector. From this upgrade, we expect an improvement in sensitivity by roughly an order of magnitude \\cite{IceCube:ICRC2013_slowMPs}. This gain is supplemented by  improvements of the data selection which have been developed after completion of this analysis. Examples are the implementation of a Kalman-filter based  HLC hit selection, which  improves the angular and speed reconstruction, and the implementation of an event selection based on a Boosted Decision Tree \\cite{Zierke:Masterarbeit}."
    },
    "1402/1402.4130_arXiv.txt": {
        "abstract": "The recently discovered dwarf galaxy Perseus I appears to be associated with the dominant plane of non-satellite galaxies in the Local Group (LG). We predict its velocity dispersion and those of the other isolated dSphs Cetus and Tucana to be 6.5, 8.2, and $5.5\\,\\mathrm{km\\,s}^{-1}$, respectively. The NGC 3109 association, including the recently discovered dwarf galaxy Leo P, aligns with the dwarf galaxy structures in the LG such that all known nearby non-satellite galaxies in the northern Galactic hemisphere lie in a common thin plane (rms height 53 kpc; diameter 1.2 Mpc). This plane has an orientation similar to the preferred orbital plane of the Milky Way (MW) satellites in the vast polar structure. Five of seven of these northern galaxies were identified as possible backsplash objects, even though only about one is expected from cosmological simulations. This may pose a problem, or instead the search for local backsplash galaxies might be identifying ancient tidal dwarf galaxies expelled in a past major galaxy encounter. The NGC 3109 association supports the notion that material preferentially falls towards the MW from the Galactic south and recedes towards the north, as if the MW were moving through a stream of dwarf galaxies. ",
        "introduction": "The Milky Way (MW) is surrounded by a vast polar structure (VPOS) of satellite objects including the satellite galaxies, young halo globular clusters and several stellar and gaseous streams \\citep*{1976MNRAS.174..695L,2012MNRAS.423.1109P}. The proper motions of the 11 classical satellite galaxies reveal that these almost exclusively co-orbit in this VPOS, which allowed us to predict the proper motions of the remaining satellite galaxies \\citep{2013MNRAS.435.2116P}.  \\citet*{2005A&A...431..517K} have first identified this planar alignment as being inconsistent with cosmological simulations based on the cold dark matter paradigm with a cosmological constant, $\\Lambda$CDM. This finding subsequently triggered an ongoing debate on whether such structures can be reconciled with cosmological expectations \\citep[e.g.][]{2008ApJ...686L..61D,2008MNRAS.385.1365L,2009ApJ...697..269M,2009MNRAS.399..550L,2011MNRAS.415.2607D,2012MNRAS.424...80P,2013MNRAS.429.1502W,2013MNRAS.435.2116P} or rather points at a different origin such as the formation of phase-space correlated tidal dwarf galaxies \\citep[TDGs, e.g.][]{2005PASJ...57..429S,2007MNRAS.376..387M,2010ApJ...725L..24Y,2011A&A...532A.118P,2012MNRAS.427.1769F,2013MNRAS.431.3543H,2013MNRAS.429.1858D}.  \\citet{2013Natur.493...62I} and \\citet{2013ApJ...766..120C} have recently discovered a similar 'Great Plane of Andromeda' (GPoA), a co-orbiting alignment consisting of about half of the satellite galaxies of the Andromeda galaxy (M31), the other major galaxy in the Local Group (LG). Motivated by this discovery that satellite galaxies appear to preferentially live in phase-space correlated structures, \\citet{2013MNRAS.435.1928P} set out to search for similar structures on a LG scale. They have discovered that all but one of the 15 LG dwarf galaxies more distant than 300\\,kpc from the MW and M31 are confined to two narrow (short-to-long axis ratios of $\\approx 0.1$) and highly symmetric planes, termed LGP1 and LGP2. LGP1 is the dominant plane both by number of objects (about nine), and alignment with additional features, such as the Magellanic Stream which traces the positions and line-of-sight (LOS) velocities of the LGP1 plane members in the southern Galactic hemisphere.  Given that the number of known dwarf galaxies in the LG more distant than 300\\,kpc from both major galaxies is still low, each additional detection poses a chance to test the existence of the planar LG structures and to potentially refine our understanding of these structures. Such an opportunity is now provided by the recent discovery of the dwarf spheroidal (dSph) galaxy Perseus I at a distance of 374\\,kpc from M31 \\citep{2013arXiv1310.4170M}. In the following, we test whether it can be considered to be associated to either LGP1 or LGP2.  In addition, we predict the velocity dispersion of Perseus I and two other non-satellite dSphs as expected in Modified Newtonian Dynamics \\citep[MOND,][]{1983ApJ...270..365M,2012LRR....15...10F}. Similar predictions have been made for other dwarf galaxies in M31's vicinity \\citep{2013ApJ...766...22M} and these have successfully passed the test of observations \\citep{2013ApJ...775..139M}. Unfortunately, no similar predictions are possible in the $\\Lambda$CDM framework.  Another recently discovered nearby dwarf galaxy, Leo P \\citep{2013AJ....146...15G,2013AJ....145..149R}, has lead \\citet{2013arXiv1310.6365B} to re-investigate the NGC 3109 association, a group of dwarf galaxies at a distance of about 1.3--1.4\\,Mpc from the MW that consists of NGC 3109, Antlia, Sextans A and Sextans B \\citep{1999ApJ...517L..97V,2006AJ....132..729T}. They realised that Leo P aligns with the four other members of the association in a very narrow, linear structure. As the NGC 3109 association is very close to the LG and has a linear extent of 1.2\\,Mpc, similar to its distance from the MW, we will investigate its orientation in the context of the LG planes of non-satellite dwarf galaxies. This reveals an intriguing alignment with the other three nearby non-satellite galaxies in the northern hemisphere of the MW and leads us to discuss suggested origins for the NGC 3109 association in light of the geometry of the LG.  The paper is structured as follows. In Sect. \\ref{sect:perseus} we determine whether Perseus I is associated with one of the dwarf galaxy planes in the LG. In Sect. \\ref{sect:MOND} we predict the velocity dispersion of the distant dSphs in the LG, Perseus I, Cetus and Tucana. In Sect. \\ref{sect:NGC3109association} we determine the orientation of the NGC 3109 association in the same coordinate system used in \\citet{2013MNRAS.435.1928P}, discuss possible origins for the alignment and conclude that the association is likely part of the LG dwarf galaxy structures. In Sect. \\ref{sect:backsplash} we discuss how the search for cosmological backsplash galaxies in the LG might give rise to two additional small-scale problems of cosmology and how it could falsely identify TDGs as backsplash objects. In Sect. \\ref{sect:scheme} we present a sketch of the LG dwarf galaxy structures and their preferred direction of motion and discuss open questions and limitations in Sect. \\ref{sect:openquestions}. Finally, we summarize our results in Sect. \\ref{sect:conclusion}. ",
        "conclusions": "\\label{sect:conclusion}  We have investigated how the recently discovered dSph galaxy Perseus I and the NGC 3109 association extended with the recently discovered dwarf galaxy Leo P are related to, and might fit in with, the dwarf galaxy structures present in the LG. Our work has shown that the NGC 3109 association cannot necessarily be interpreted as an independent group of galaxies, but might be related to the LG dwarf galaxy population and as such might provide important constraints on attempts to model the whole LG. The main results of our analysis are:  \\begin{enumerate} \\item Perseus I is consistent with being part of the LGP1$^{\\mathrm{mod}}$, the dominant plane of non-satellite galaxies in the LG, at least if Andromeda XVI is associated with the GPoA. \\item In the context of MOND, we have predicted Perseus I's velocity dispersion to be $\\sigma = 6.5^{+1.2}_{-1.0} \\pm 1.1\\,\\mathrm{km\\,s}^{-1}$. The corresponding prediction for Cetus ($\\sigma = 8.2^{+1.5}_{-1.3} \\pm 0.4\\,\\mathrm{km\\,s}^{-1}$) is in much better agreement with the more recent observational value \\citep[$\\sigma = 8.3 \\pm 1.0\\,\\mathrm{km\\,s}^{-1}$:][]{2014arXiv1401.1208K} than with the previous measurement \\citep[$\\sigma = 17 \\pm 2.0\\,\\mathrm{km\\,s}^{-1}$:][]{2007MNRAS.375.1364L}. The prediction for Tucana ($\\sigma = 5.5^{+1.0}_{-0.9} \\pm 0.4\\,\\mathrm{km\\,s}^{-1}$) is in conflict with the available measurement \\citep[$\\sigma = 15.8^{+4.1}_{-3.1}\\,\\mathrm{km\\,s}^{-1}$:][]{2009A&A...499..121F}. We note that the observations of \\citet{2009A&A...499..121F} lack the spectral resolution to resolve the predicted velocity dispersion. Unfortunately, no similar predictions are possible in a $\\Lambda$CDM context. \\item The orientation of the NGC 3109 association consisting of the dwarf galaxies NGC 3109, Antlia, Sextans A, Sextans B and Leo P has been determined in the same coordinate system used to review the planes of co-orbiting satellite galaxies around the MW and M31 and the symmetric larger-scale dwarf galaxy structure in the LG \\citep{2013MNRAS.435.1928P}. The association aligns with the Supergalactic Plane, is almost perpendicular to our line-of-sight and parallel but offset by 300-500\\,kpc to LGP1$^{\\mathrm{mod}}$. \\item The members of the NGC 3109 association have large receding velocities which indicate that they have been close to the MW in the past, possibly at the same time about 7--9\\,Gyr ago. This is consistent with their orbits passing within $\\approx 25$\\,kpc of the MW suggested by \\citet{2013MNRAS.tmp.2500S} and the identification as likely backsplash galaxies by \\citet{2012MNRAS.426.1808T}. Together with the association's extremely narrow extent and perpendicular orientation this argues against the association tracing a thin and cold cosmological filament. The timing is consistent with independent timing estimates for several suggested major galaxy encounter scenarios in the LG, during which phase-space correlated populations of TDGs could have formed that would today give rise to the observed dwarf galaxy structures \\citep{2012MNRAS.423.1109P,2013MNRAS.431.3543H,2013A&A...557L...3Z}. \\item The association aligns with the other three distant ($> 300$\\,kpc) LG galaxies in the northern hemisphere of the MW in a narrow plane (RMS height of $\\Delta\\,=\\,53$\\,kpc and axis ratios of $c/a\\,=\\,0.10$\\ and $b/a\\,=\\,0.76$). This ``Great Northern Plane'' passes through the center of the LG, is inclined to LGP1$^{\\mathrm{mod}}$\\ by only $27^\\circ$ and to the GPoA by $40^\\circ$ and is consistent with being aligned with the preferred orbital plane of the MW satellites in the VPOS. \\item Five out of seven (6 of 8 if the later discovered Leo P would be included) of the galaxies in the Great Northern Plane have been identified as likely backsplash objects by \\citet{2012MNRAS.426.1808T}, and most of the remaining five backsplash candidates in the southern Galactic hemisphere are also situated close to the same plane. As only a small fraction of sub-haloes in simulations are identified as backsplash objects the finding of a majority of such galaxies in one hemisphere is extremely unlikely ($\\approx 0.1$\\ per cent) and might constitute an ``overabundant backsplash problem'' for $\\Lambda$CDM. It would mean that more galaxies are receding in one direction from the MW than are being accreted onto the MW from that direction. A natural explanation for this would be the local formation of galaxies as TDGs, which, if expelled to large distances, have very similar orbital properties like cosmological backsplash galaxies. That the backsplash candidates preferentially lie in a common plane is also consistent with an interpretation as phase-space correlated TDGs. \\item LG galaxies are found to be preferentially infalling in the Galactic south and receding in the Galactic north, which possibly indicates that the MW is moving through a stream of dwarf galaxies. This would be in qualitative agreement with the M31-merger scenario by \\citet{2010ApJ...725..542H}, in which our Galaxy passes through the tidal debris expelled in a past merger forming M31 \\citep[see also][]{2010ApJ...725L..24Y,2012MNRAS.427.1769F,2013MNRAS.431.3543H}. However, the TDG origin would be in M31, such that the galaxies move away from M31 sufficiently fast which would imply that they have a tangential velocity component relative to the MW, rendering the timing estimate of Sect. \\ref{sect:timing} somewhat useless. The scenario should however be testable on LG scales if proper motion measurements of the distant dwarf galaxies could be obtained. \\end{enumerate}"
    },
    "1402/1402.3474_arXiv.txt": {
        "abstract": "{To deepen our understanding of the role of small-scale magnetic fields in active regions (ARs) and in the quiet Sun (QS) on the solar irradiance, it is fundamental to investigate the physical processes underlying their continuum brightness. Previous results showed that magnetic elements in the QS reach larger continuum intensities than in ARs at disk center, but left this difference unexplained.} {We use Hinode/SP disk center data to study the influence of the local amount of magnetic flux on the vigour of the convective flows and the continuum intensity contrasts.} {The apparent (i.e. averaged over a pixel) longitudinal field strength and line-of-sight (LOS) plasma velocity were retrieved by means of Milne-Eddington inversions (VFISV code). We analyzed a series of boxes taken over AR plages and the QS, to determine how the continuum intensity contrast of magnetic elements, the amplitude of the vertical flows and the box-averaged contrast were affected by the mean longitudinal field strength in the box (which scales with the total unsigned flux in the box).} {Both the continuum brightness of the magnetic elements and the dispersion of the LOS velocities anti-correlate with the mean longitudinal field strength. This can be attributed to the ``magnetic patches'' (here defined as areas where the longitudinal field strength is above 100 G) carrying most of the flux in the boxes. There the velocity amplitude and the spatial scale of convection are reduced. Due to this hampered convective transport, these patches appear darker than their surroundings. Consequently, the average brightness of a box decreases as the the patches occupy a larger fraction of it and the amount of embedded flux thereby increases.} {Our results suggest that as the magnetic flux increases locally (e.g. from weak network to strong plage), the heating of the magnetic elements is reduced by the intermediate of a more suppressed convective energy transport within the larger and stronger magnetic patches. This, together with the known presence of larger magnetic features, could explain the previously found lower contrasts of the brightest magnetic elements in ARs compared to the QS. The inhibition of convection also affects the average continuum brightness of a photospheric region, so that at disk center, an area of photosphere in strong network or plage appears darker than a purely quiet one. This is qualitatively consistent with the predictions of 3D MHD simulations.} ",
        "introduction": "\\label{sec_intro}  Variations of the Total Solar Irradiance (TSI) on timescales from days to the solar cycle can be reasonably well reproduced based on the evolving distribution of solar surface magnetic fields \\citep[which is the basis of e.g. the SATIRE reconstructions,][]{Fligge00, Solanki02, Krivova03, Wenzler06}. Although the TSI reconstructions proved very successful \\citep[with a correlation of up to 0.98 with the measured irradiance,][]{Ball11}, they are based on an empirical relationship between the emergent intensity of a pixel and the measured magnetic field (via MDI magnetograms) whose simple linearity lacks physical basis; its justification relies on its ability to reproduce the disk-average irradiance and the convenience of using a single model atmosphere for all faculae \\citep[for more details see][]{Fligge00}. In particular, the spectral radiance of a given pixel is assumed to be \\emph{uniquely} determined by the magnetic signal at that pixel, thus neglecting its magnetic environment and location on the Sun.  The surface magnetic field outside Sunspots is mainly distributed in active region (AR) plages and in the quiet Sun (QS) network outlining supergranular cells. There the magnetic flux is often concentrated into features traditionally described in terms of flux tubes, with field strengths of the order of kG \\citep[e.g.][]{Frazier72, Stenflo73, Rabin92a, Ruedi92} and a spectrum of sizes \\citep[see][for a review]{Solanki_rev06}. At the lower end of this size spectrum are found the so-called ``magnetic elements'' \\citep[see][]{Schuessler_rev92, Solanki_rev93}, which have been spatially resolved recently by the balloon-borne SUNRISE ImaX observations at an angular resolution of $\\sim 0\\dotarsec14$ \\citep[][]{Lagg10}. As a result of being hotter than their surroundings at equal optical depth \\citep{Schuessler_rev92}, they often appear brighter than the average quiet photosphere, i.e. they have a positive ``contrast''. This is particularly pronounced when observing in the core of spectral lines and in molecular bands \\citep{Chapman68, Sheeley69, Muller84}, but also holds at continuum wavelengths or broader wavelength bands, even at disk center, if the magnetic elements are sufficiently resolved \\citep[see][the two last citations refering to radiation-MHD simulations]{Muller83b, Foukal84, Lawrence88, Riethmueller10, Schuessler88, Voegler03}. Note that the contrast of magnetic features further increases from the disk center to the limb in a way that is still a matter of debate \\citep{Steiner07} and shall be treated in a forthcoming paper. Therefore, taking into account the larger area coverage of magnetic elements compared to sunspots, their radiance overcompensates the solar brightness deficit due to sunspots at activity maximum \\citep{Froehlich00}. It is known that ARs typically harbour more magnetic flux per unit area than the network and thus a higher number density of magnetic elements. The larger available flux in ARs also leads to the formation of larger and darker magnetic features such as pores and micropores, the latter being small (sub-arcsec) magnetic features that are darker than the QS continuum \\citep[][]{Beckers68, Tarbell77}. The inner part of supergranulation cells, the so-called ``internetwork'' (IN), also contains small magnetic elements but in lower number density \\citep[][]{Muller83a, Lites02, Dominguez03, deWijn05, deWijn08}. Finally, the solar photosphere contains weaker equipartition fields everywhere \\citep[detected mostly in the IN by, e.g.][]{Lin95, Solanki96, Lites02, Khomenko03}, but the latter have been estimated to bring negligible contributions to the TSI variations \\citep{Schnerr10}.  To determine how these different components of solar surface magnetism contribute to the solar irradiance, one should investigate the intensity-magnetic field relation within ARs, the QS network and even the IN. In the 1990's, \\citet{Title89, Topka92, Lawrence93} and \\citet{Topka97} performed a series of ground-based studies of continuum intensity vs. magnetogram signal in ARs and in the QS at disk center, and found that the QS continuum intensity contrast reaches larger values than in ARs at equal magnetogram signal. This early result was confirmed by our recent study using Hinode/SP data at a constant and higher spatial resolution than the earlier studies \\citep[][hereafter Paper I]{Kobel11a}. Note that in our study, the contrast of magnetic elements could be interpreted as a measure of their intrinsic brightness by virtue of the comparable contrast references used in ARs and in the QS. We also found that the bright magnetic elements in ARs and in the QS share similar filling factors (along with similar kG field strengths and inclinations close to vertical). Assuming that, at Hinode's resolution and at disk center, the filling factor reflects the size of the unresolved magnetic elements (at the height of line formation), this questions the conventional interpretation that the brightness of magnetic features is \\emph{primarily} dictated by their size \\citep[][see also the Introduction of Paper I]{Spruit81}.  As discussed in Paper I, the other factor that could possibly influence the brightness of magnetic elements (besides their size) is the surrounding convective energy transport, because it determines the energy available to be radiated into the flux tubes through their ``hot walls'' \\citep[][]{Spruit76}. Since AR plages typically have a larger mean field strength than the QS network, the Lorentz force should inhibit the convective flows more strongly in ARs. The granulation indeed appears ``abnormal'' in plages \\citep[e.g.][]{Dunn73}, such that its spatial scale is reduced compared to the QS, as are both the vertical and horizontal components of the flow field \\citep{Schmidt88, Title89, Title92, Keller91}. 3D MHD simulations also disclose a similar behaviour as the mean field increases \\citep{Voegler05b}. Furthermore, these carefully set up simulations (with constant inflowing entropy density at the bottom boundary) demonstrated that the convective energy transport was increasingly inhibited to the point of altering the \\emph{vertical} radiative energy output from the boxes, thus providing an explanation for magnetic elements appearing darker in more active environments at disk center.  The work presented herein is an attempt to explain the effect of the local amount of magnetic flux on the continuum brightness of magnetic elements and on the local average continuum brightness of the solar surface at disk center. Inspired by MHD simulations, we carried out a ``local box analysis'' that considers a series of small square fields of view taken from different regions of QS and ARs (see Sect. \\ref{sec_boxes}) located at disk center, observed with the spectropolarimeter instrument onboard Hinode. In Sect. \\ref{sec_results}, we examine the relation between the contrast and the longitudinal field strength inside these boxes. In particular, we assess the influence of the mean longitudinal field strength in the box on the intensity contrast of the magnetic elements, on the vigour of the vertical convective flows, and on the contrast averaged over the box. One reason to restrict our study to disk center is that in this way we can interpret the Doppler velocities retrieved by Milne-Eddington inversions from the data (see Sect.\\ref{sec_scans_inversions}) in terms of vertical flows. These results are then discussed in Sect. \\ref{sec_disc}, with special emphasis on the role of the magnetically-suppressed convection on the contrasts. ",
        "conclusions": "\\label{sec_concl}  Previous studies found that at disk center, the continuum intensity contrast in the QS network reached higher values than in ARs \\citep{Title89, Topka92, Lawrence93} and \\citep{Topka97}. This was confirmed by the first paper of this series \\citep[][Paper I]{Kobel11a} using Hinode/SP data, where we established that the relation between continuum brightness (at 630.2 nm) and longitudinal field strength $B_{\\rm app,los}$ (retrieved by Milne-Eddington inversions) exhibits a higher peak in the QS network than in ARs. Since the magnetic elements producing the peak share similar magnetic filling factors, we argued that this brightness difference between the QS and ARs is unlikely to be explained solely by a difference in the size of the brightest magnetic elements (although the presence of larger features in ARs certainly plays a role in determining their brightness more generally, see below).  In this paper, we tested whether the contrast of the magnetic elements at disk center could be affected by an altered convective transport due to the local concentration of magnetic flux. With this aim, we extracted from Hinode/SP scans a series of ``boxes'' containing different amounts of magnetic flux, covering areas from AR plages to weak QS network. We found that the contrast of the brightest magnetic elements (i.e. the peak of the contrast vs. $B_{\\rm app,los}$) continuously decreases with the mean longitudinal field strength $\\left< B_{\\rm app,los} \\right>$ in the boxes (Fig. \\ref{fig_CVlos_vs_avgB}a). Since the $\\left< B_{\\rm app,los} \\right>$ is generally larger in ARs than in the QS, this can be taken as a generalization of the earlier finding that magnetic elements in the network are brighter than in ARs.  Hence, as the amount of flux (and thereby $\\left< B_{\\rm app,los} \\right>$) increases in a local area, the observed contrast of the magnetic elements associated to the contrast peak decreases due to a combination of two effects: (i) The known presence of larger (and thus darker) magnetic features, even if pores are removed. These can either provide pixels. Such larger magnetic features reduce the contrast of nearby magnetic elements through the telescope point spread function. Yet this effect is likely to play only a minor role in the decrease of the magnetic element contrasts with $\\left< B_{\\rm app,los} \\right>$ (see discussion in Sect. \\ref{sec_disccontrast}). (ii) The disturbance of convective energy transport in the surrounding of the magnetic features. This is supported by the anti-correlation between the fluctuation of the longitudinal velocities outside the magnetic features and $\\left< B_{\\rm app,los} \\right>$ (Fig. \\ref{fig_CVlos_vs_avgB}b) and the lower brightness of the abnormal granulation surrounding magnetic elements compared to weakly magnetized areas (Fig. \\ref{fig_avgC_vs_avgB_patches}b).  A closer inspection of the strong network and plage boxes revealed that the magnetic flux in these boxes is mostly carried by ``magnetic patches'' that are well delimited by $B_{\\rm app,los} > 100$ G (Fig. \\ref{fig_Phi_fA_vs_avgB}a). The flux density in these patches is such that the vertical convective velocities inside them (around magnetic features) are reduced significantly relative to their surroundings (the rms($v_{\\rm los}$) is about 100 m s$^{-1}$ lower), this reduction being more important as the mean field of the patches increases (cf. Fig. \\ref{fig_rmsVlos_vs_avgB_patches}b). These reduced velocities probably limit the heat upflow and consequently lower the brightness of the granulation in the patches compared to the granulation in the weakly magnetized areas (Fig. \\ref{fig_avgC_vs_avgB_patches}b). The reduced spatial scale of the (abnormal) granulation inside the patches probably also contributes to decreasing the efficiency of the convective energy transport. Even with their magnetic elements, the patches are generally darker than their surroundings (Fig. \\ref{fig_avgC_vs_avgB_patches}a), and become all the more dark when $\\left< B_{\\rm app,los} \\right>$ increases in the box and the patches occupy a larger part of it (Fig. \\ref{fig_Phi_fA_vs_avgB}b). We interpret the patches as having the following influence on the brightness of the magnetic elements: \\emph{the more extended the magnetic patch is (independently of the box size) and the larger its mean field, the larger is the inhibition of convection, resulting in cooler granules and smaller heat influx to the ``hot walls'' of the magnetic elements which become cooler and less bright (at least at disk center)}.  Since the patches are darker than their surrounding and occupy a larger part of the box as $\\left< B_{\\rm app,los} \\right>$ increases, the average brightness of a box at disk center also decreases with increasing $\\left< B_{\\rm app,los} \\right>$ (by about 1.2\\% from weak network to plage boxes). Previous 3D MHD simulations of \\citet{Voegler05b} already predicted a reduction of the vertical bolometric intensity of boxes with mean field strength $\\left< B\\right> > 200$ G. The qualitative agreement between our results and these simulations (for $\\left< B\\right> > 200$ G) supports their use to model the different component atmospheres used in irradiance reconstructions \\citep[see, e.g.,][]{Unruh09}.  Up to now, the effect of the inhibition of convection on the photospheric contrast has not been taken into account in irradiance reconstruction models. These models only use the photometric area of magnetic features \\citep{Penza03, Fontenla05} or their ``filling factor'' within a pixel \\citep[in the case of SATIRE, see][for reviews]{Solanki05, Krivova10} as the time-dependent variables modulating the intensity output given by semi-empirical 1D model atmospheres. Our results suggest that the reconstructions could be possibly enhanced by modulating the emitted intensity (e.g. for a given magnetic filling factor) according to the surrounding magnetic flux in a local (e.g. in a $20\\arsec \\times 20\\arsec$) area. However, the finding that the inhibition of convection in more active areas affects the brightness of magnetic elements has been tested here for the disk center only. In the next article of this series, we shall investigate this effect away from the disk center by comparing center-to-limb variation of the continuum contrast of magnetic elements in the QS and in ARs separately."
    },
    "1402/1402.4853_arXiv.txt": {
        "abstract": "X--ray photons coming from an X--ray point source not only arrive at the detector directly, but also can be strongly forward-scattered by the interstellar dust along the line of sight (LOS), leading to a detectable diffuse halo around the X--ray point source. The geometry of small angle X--ray scattering is straightforward, namely, the scattered photons travel longer paths and thus arrive later than the unscattered ones; thus the delay time of \\xshalo\\ photons can reveal information of the distances of the interstellar dust and the point source. Here we present a study of the \\xshalo\\ around \\src , which is one of the so--called supergiant fast X--ray transients. \\src\\ underwent a striking outburst when observed with \\chd\\ on 2004 July 3, providing a near $\\delta$--function lightcurve. We find that the X--ray scattered halo around \\src\\ is produced by two interstellar dust clouds along the LOS. The one which is closer to the observer gives the \\xshalo\\ at larger observational angles; whereas the farther one, which is in the vicinity of the point source, explains the halo with a smaller angular size. By comparing the observational angle of the scattered halo photons with that predicted by different dust grain models, we are able to determine the normalized dust distance. With the delay times of the scattered halo photons, we can determine the point source distance, given a dust grain model. Alternatively we can discriminate between the dust grain models, if the point source distance is known independently. ",
        "introduction": "Immersed in the interstellar medium (ISM), the interstellar dust grains not only absorb X--ray photons but also scatter X--ray photons. \\citet{ove65} first predicted the presence of \\xshalo s around X--ray sources. Nonetheless due to the limitation of the angular resolution of X--ray imaging telescopes, it was not until two decades later, \\citet{rol83} first observed the diffuse \\xshalo\\ around GX 339$-$4 with the Imaging Proportional Counter (IPC) instrument onboard \\ein . The theory of small angle X--ray scattering has since been refined by a number of authors \\citep[e.g.,][etc]{mau86,mat91,smi98}. Meanwhile, X--ray scattering phenomena around both X--ray point sources and extended supernova remnants (SNRs) were found with \\ros\\ \\citep[e.g.,][]{pre95}, \\chd\\ \\citep[e.g.,][]{smi02}, \\xmm\\ and \\swi /XRT \\citep[e.g.,][]{tie10}, etc.  Likewise, albeit it had already been pointed out by \\citet{tru73} that the time delay effect in the \\xshalo s can be used to determine the distances of variable X--ray point sources, it was nearly three decades later then came its long overdue application. \\citet{pre00} roughly estimated the distance of Cyg X$-$3, which was observed with the Advanced CCD Imaging Spectroscopy (ACIS) instrument onboard \\chd . Thanks to the fine angular resolution of \\chd , \\citet{tho09},~\\citet{lin09}, and~\\citet{xia11} used the time delay effect in the \\xshalo s to determine the distances of Cen X$-$3, Cyg X$-$3, and Cyg X$-$1, respectively. Additionally, \\citet{vau04} first expanded the scope onto $\\gamma$--ray bursts (GRBs), in which case an even simpler geometry was involved. The evolving \\xsring s around GRB~031203 \\citep{vau04, tie06}, GRB~050724 \\citep{vau06}, GRB~050713A \\citep{tie06}, GRB~061019 and GRB~070129 \\citep{via07} were detected with both \\xmm\\ and \\swi /XRT.  It is indisputable that the characteristics of the X--ray dust scattered halos (e.g. the radial profile, etc) depend upon the properties of interstellar dust grains, including chemical composition, dust size distribution, dust spatial distribution, and sometimes the morphology and the alignment of elongated dust grains might play important roles. So far, various types of interstellar dust grain models have been established based on different observational results such as interstellar extinction, diffuse infrared emission feature and so forth \\citep[see][for a review]{lia03}. Among them, the three types of models provided by \\citet[][hereafter MRN]{mat77}, \\citet[][hereafter WD01]{wei01}, \\citet[][hereafter ZDA]{zub04} are most widely used.  In this work we analyze the archival \\chd\\ data of \\src , which was first discovered with \\intg\\ on 2003~September~17 \\citep{sun03} as a galactic high mass X--ray binary (HMXB). With several subsequent observations carried out with \\xmm\\ \\citep{gri04}, \\chd\\ \\citep{iza05}, \\swi /XRT \\citep{sid09}, and \\suz\\ \\citep{ram09}, remarkable outbursts with a duration time scale of hours were spotted. Along with the confirmation of an O9Ib blue supergiant donor in the binary system \\citep{pel06}, \\src\\ was confirmed as one of the so--called supergiant fast X--ray transients \\citep[SFXTs][]{sgu05,smi06,neg06}. SFXT, a subclass of HMXB, is characterised by the presence of a supergiant companion and significant outbursts lasting typically a few hours. Typically, the peak luminosity of a flare can be about a factor of 10$^{3}$--10$^{5}$ times the fainter quiescent X--ray luminosity.  In Section~\\ref{sct:data}, we analyze the \\chd\\ ACIS-S data of \\src , focusing on timing (Section~\\ref{sct:timing}) and imaging analysis (Section~\\ref{sct:imaging}). Then, we derive the time lags of the \\xshalo\\ photons via cross correlation method in Section~\\ref{sct:ccf}. In Section~\\ref{sct:theta}, we model the deviation of the arithmetic mean of the observation angle from the mid-value of the angular distance of the annular region. We subsequently present the distance measurement for interstellar dust clouds along the line of sight (LOS) and the point source (Section ~\\ref{sct:dist}). A dynamical distance measurement for \\src\\, obtained from a Galactic Center molecular clouds survey, presented in Section \\ref{sct:clouds_dist}, is consistent with the estimated point source distance. In Section~\\ref{sct:caveats}, we briefly discuss the feasibility of a promising application of the relationship between interstellar dust grain models and the average observational angles for annular halos. Finally, we summarize our results in Section \\ref{sct:summary}. ",
        "conclusions": "\\label{sct:discussion} \\subsection{Kinematic distance measurements of molecular clouds} \\label{sct:clouds_dist} For $d\\sim3.6$~kpc \\citep{rah08}, the distances for the closer and farther dusts are $\\sim1.8$~kpc and $\\sim3.4$~kpc away from us, respectively. In this subsection we try to find kinematic distance measurements of the molecular clouds along the LOS toward \\src\\, for comparison with our geometrical distances of the dust slabs. A rough estimate of the radial velocity of the molecular clouds where the dust slab is embedded can be made via \\citep{rdu09} \\begin{equation} r=R_\\odot \\sin l \\frac{V(r)}{v_{\\rm los}+V_\\odot \\sin l}, \\label{v_los} \\end{equation} where $r$ is the distance of the molecular cloud to the Galactic Center (GC), $R_{\\odot}$ is the distance of the Sun to GC, {\\it l} is the galactic longitude of the LOS, $V_{\\odot}$ is the rotation velocity of the Sun, $V(r)$ is the rotation velocity of the molecular cloud, and $v_{\\rm los}$ is the projection of $V(r)$ to the LOS, also known as the radial velocity. Assuming $R_{\\odot}=8.5$~kpc, $V_{\\odot}=220$~km s$^{-1}$ and a flat rotation curve (i.e. $V(r)=220$~km s$^{-1}$), we derive $v_{\\rm los}$ of $\\sim3.1$~km s$^{-1}$ and $\\sim7.8$~km s$^{-1}$ for the closer and farther molecular cloud, respectively. In fact, \\citet{dah98} reported an averaged $v_{\\rm los}\\sim5$~km s$^{-1}$ for $^{12}$CO(1-0) emission line for the Southern Clump 2 region, $l\\in(2\\fdg7, 3\\fdg5)$, $b\\in(0\\fdg15, 0\\fdg35)$, which is roughly consistent with our result of the farther dust.  \\subsection{Issues with observational angles} \\label{sct:caveats} In Section~\\ref{sct:theta}, we have shown that given the normalized distance of a dust slab, different interstellar dust models predict different average observational angles for scattered photons within certain annular regions. This offers the advantage of estimating the parameter $x$ from data, thus breaking the degeneracy between $x$ and $d$. We tested several interstellar dust grain models with the data of the farther dust, for which the predictions of these models become quite different, because the scattering angle $\\theta_{\\rm sca}=\\theta_{\\rm obs}/(1-x)$ is quite large for the farther dust. It is therefore possible to distinguish among different interstellar dust models, as demonstrated above. However in the calculations of the scattering differential cross section, Gaussian approximation is used for the form factor in the RG approximation. \\citet{smi98} pointed out that the Gaussian approximation leads to deviations at large scattering angle ($\\gtrsim200~{\\rm arcsec}$) and large dust grain size. In our case, i.e., $x>0.95$ and $\\theta_{\\rm obs}\\sim10~{\\rm arcsec}$ for the farther dust slab, we have $\\theta_{\\rm sca}>200$~arcsec. Moreover, the upper limits of the grain size are greater than 0.3~micron, and even reach to 0.9~micron. Therefore the Gaussian approximation may cause considerable inaccuracies to the model predictions. Alternatively, one can turn to Mie theory \\citep{vhu57}, which is more accurate for $E\\sim1~$keV and/or large scattering angles, but numerically more difficult to carry out the calculations.   In this work, we analyzed the \\xshalo\\ around \\src\\ with cross correlation method. The main results are summarized as follows: \\begin{enumerate} \\item From the cross correlation functions between the streak lightcurve (used as a proxy for the point source lightcurve) and the lightcurves of the annular halos, we conclude that there are at two interstellar dust clouds along the LOS toward \\src .  \\item By comparing the observational angle of the scattered halo photons with that predicted by different dust grain models, the normalized dust distance can be determined independent of the distance of the point source.  \\item Given the point source distance of $\\sim3.6$~kpc, the closer dust, which is $\\sim 1.8$~kpc away from us, is responsible for \\xshalo s at larger observational angles ($\\gtrsim12.5$ arcsec). The farther dust, which is quite close to the point source ($\\sim3.4$~kpc away from us), explains the X--ray scattered halos at smaller angular distances ($\\lesssim12.5$ arcsec).  \\item We determined the model-dependent point source distances, which are compared with that yielded by IR observations. We find that among the 18 tested dust grain models (MRN, WD01, ZDAs and XLNW), the four dust grain models COMP-AC-S/B, COMP-NC-S/FG are better, COMP-GR-S/FG, COMP-AC-FG, COMP-NC-B are also acceptable, but the rest dust grain models fail to obtain consistent source distance. \\end{enumerate}  Similar to the GRBs, the transient nature of SFXTs can, in principle, be used to precisely determine the geometrical distances of interstellar dust and the point source by taking advantage of the time delay effect of the small angle X--ray scattering phenomena. However, we have to face some practical difficulties, such as: 1) the angular resolutions of the space telescopes are relatively poor; 2) the effective area is small and thus the photon counts are relatively low; or 3) for observations of \\xshalo\\ caused by dust slab in the vicinity of the point source, the time lags can be quite large, but no observations with sufficiently long effective exposure times are available. The effective exposure time refers to the exposure time for observing the \\xshalo . For instance, in our work, since the outburst occurs $\\sim$~10~ks after the beginning of this observation, the effective exposure time for observing the \\xshalo\\ here is only $\\sim$~9 ks.  With the fine angular resolution, \\chd\\ has the capability to observe \\xshalo s around SFXTs especially at smaller angular distance, although the collecting area of ACIS is small. Unfortunately, insofar as the archival \\chd\\ data, only \\src\\ (ObsID 4550) allows us to study the \\xshalo . In terms of other observations of SFXTs, either the exposure time is only several kiloseconds (e.g. XTE J1739$-302$), or no flaring activity was caught (e.g., for IGR J19410$-$0951). Therefore, we suggest that in the future more long term follow up observations of the outbursts of SFXTs shall be made with \\chd\\ to study the \\xshalo\\ and thus the interstellar dust models in further details."
    },
    "1402/1402.1192_arXiv.txt": {
        "abstract": "We have determined spectroscopic orbits for five single-lined spectroscopic binaries, HD 100167, HD 135991, HD 140667, HD 158222, HD 217924. Their periods range from 60.6 to 2403 days and the eccentricities, from 0.20 to 0.84. Our spectral classes for the stars confirm that they are of solar type, F9 to G5, and all are dwarfs. Their [Fe/H] abundances, determined spectroscopically, are close to the solar value and on average are 0.12 greater than abundances from a photometric calibration.  Four of the five stars are rotating faster than their predicted pseudosynchronous rotational velocities. ",
        "introduction": "Nearly a decade ago a group of solar-type stars with $V$ magnitudes between 5.9 and 8.1 was selected to expand a photometric program that was designed to detect long-term spot cycles in dwarfs similar to the Sun \\citep{h99}.  Although most of these moderately bright stars had radial velocities published in the literature, many did not.  To remove binaries from the photometric sample, we obtained spectra and measured radial velocities.  Our spectroscopic observations soon identified five single-lined binaries, HD~100167, HD~135991, HD~140667, HD~158222, and HD~217924, which we continued to observe.  The basic properties of these binaries are given in Table~\\ref{tbl-basic}.  During the course of our extended spectroscopic observing campaign, others have also identified these five stars as binaries.  From our observations we determine spectroscopic orbits and spectral and luminosity classes, measure iron abundances and $v$~sin~$i$ values, and then briefly discuss the systems. ",
        "conclusions": ""
    },
    "1402/1402.5778_arXiv.txt": {
        "abstract": "{Transition disks are protoplanetary disks with inner depleted dust cavities that are excellent candidates for investigating the dust evolution when there is a pressure bump. A pressure bump at the outer edge of the cavity allows  dust grains from the outer regions to stop their rapid inward migration towards the star and to efficiently grow to millimetre sizes. Dynamical interactions with planet(s) have been one of the most exciting theories to explain the clearing of the inner disk.} {We look for evidence of  millimetre dust particles in transition disks by measuring their spectral index  $\\alpha_{\\mathrm{mm}}$ with new and available photometric data. We investigate the influence of the size of the dust depleted cavity on the disk integrated millimetre spectral index.} {We present the 3-millimetre (100~GHz) photometric observations carried out with the Plateau de Bure interferometer of four transition disks: LkH$\\alpha$~330, UX~Tau~A, LRLL~31, and LRLL~67.  We used the available values of their fluxes at 345~GHz to calculate their spectral index, as well as the spectral index for a sample of twenty transition disks. We compared the observations with two kinds of models. In the first set of models, we considered coagulation and fragmentation of dust in a disk in which a cavity is formed by a massive planet located at different positions.  The second set of models assumes disks with truncated inner parts at different radii and with power-law dust-size distributions, where the maximum size of grains is calculated considering turbulence as the source of destructive collisions.} {We show that the integrated spectral index is higher for transition disks (TD) than for regular protoplanetary disks (PD) with mean values of $\\bar{\\alpha}_{\\mathrm{mm}}^{\\mathrm{TD}}=2.70\\pm0.13$ and $\\bar{\\alpha}_{\\mathrm{mm}}^{\\mathrm{PD}}=2.20\\pm0.07$ respectively. For transition disks, the probability that the measured spectral index is positively correlated with the cavity radius is 95\\%. High angular resolution imaging of transition disks is needed to distinguish between the dust trapping scenario and the truncated disk case.} {} ",
        "introduction": "\\label{introduction}  Observations by infrared telescopes done with  e.g. Spitzer and Herschel and by millimetre observations with radio-interferometers such as the Submillimetre Array (SMA),  the Very Large Array (VLA) or the Combined Array for Research in Millimeter-wave Astronomy (CARMA) have revealed different types of structures in circumstellar disks surrounding young stars.  Transition disks show a lack of mid-infrared radiation in the spectral energy distribution (SED), implying an absence of warm dust in the inner disk \\citep[e.g.][]{calvet2005, pietu06, espaillat10}. Some of these cavities have been resolved by continuum observations \\citep[see][for a review]{williams11}.  The recent  discoveries of  companion candidates inside transition disks, such as HD142527 \\citep{biller2012} and HD169142 \\citep{quanz2013}, promote the idea that holes in transition disks are the result of a brown dwarf, planet, or multiple planets interacting with the disk. However, none of these companion candidates have been confirmed, and alternative explanations have been suggested to  explain the observations \\citep{olofsson2013}.  Hydrodynamical simulations of planet-disk interactions have been explored for  several years  \\citep[e.g.][]{lin79, paardekooper04}. The presence of a massive planet in a circumstellar disk not only affects the surrounding gas, but also the dust distribution in the disk. A planet of $\\gtrsim$~1~$M_{\\mathrm{Jup}}$ opens a gap in a viscous disk and dust can flow through a gap, depending on the particle size and the shape of the gap \\citep{rice06, pinilla12a, zhu2012, zhu2013}. The overall dust distribution is unknown, and observations at different wavelengths with high sensitivity and resolution are crucial for understanding how dust evolves under  disk clearing environments.  Maps of sub-millimetre and microwave emission from the diffuse interstellar dust have shown that the typical values for the dust opacity index $\\beta_{\\rm{mm}}$ are $\\sim~1.7-2.0$ \\citep[see e.g.][]{Finkbeiner1999}. Hence, if the dust in protoplanetary disks is similar to the interstellar medium (ISM) dust, $\\beta_{\\rm{mm}}$ is expected to be similar. However, changes in the opacity between the ISM and denser regions  have been found \\citep[e.g.][]{henning1995}. Dust evolution in protoplanetary disks, and especially the process of grain growth to millimetre sizes in the midplane of disks, has been confirmed by observations \\citep[see reviews by][]{natta2007, testi2014}.  The slope of the SED, known as the spectral index $\\alpha_{\\mathrm{mm}}$, in the millimetre regime can be interpreted in terms of grains size, via  $\\alpha_{\\mathrm{mm}} \\approx \\beta_{\\rm{mm}} +2$ \\citep[see e.g.][]{draine2006},  and low values of  $\\alpha_{\\mathrm{mm}}$ ($\\lesssim 3.5$) correspond to big particles ($\\gtrsim$~mm).  Sub- and millimetre observations of classical  protoplanetary disks have shown that the spectral index reaches low values $\\alpha_{\\mathrm{mm}}\\lesssim 3.5$ in Taurus, Ophiuchus, Orion, Chamaeleon, Lupus, etc., star forming regions \\citep[e.g.][]{testi2001, Testi2003, rodmann2006, lommen2007, lommen2009, lommen2010, Ricci2010a, Ricci2010b, Ricci2011, ubach2012}.  Alternative explanations for the low spectral indices in disks have been addressed by different authors. One possible explanation is that the emission or part of the emission may come from unresolved regions where the dust  is optically thick at millimetre wavelengths \\citep{testi2001, Testi2003}, but it has been shown that this is unlikely to be the general explanation for the low values of the opacity index measured in large samples of disks \\citep{Ricci2012b}. In addition,  the chemical composition \\citep{Semenov2003}, porosity \\citep{henning1996, stognienko1996}, and  geometry of the dust particles can affect the opacity index. However, their impact is not significant, for reasonable values of the properties of the dust grains  \\citep[see e.g.][]{draine2006}. Thus, the most likely explanation is that dust grains have grown to mm-sizes \\citep[see in addition][]{calvet2002, natta2004a, natta2007}.  One of the most critical challenges of  planetesimal formation is to understand how the rapid inward migration of dust grains can stop. The theory of radial drift and observations at mm-wavelength are in strong disagreement \\citep[e.g.][]{brauer07}. The confirmation of long-lived presence of mm grains with observations made theoreticians to look for the  missing piece of this rapid inward-drift puzzle. Particle trapping in pressure maxima (i.e.\\ locations where $\\nabla p=0$ and $\\nabla^2p<0$) is a result of the dynamic phenomenon that, due to friction with the gas, particles tend to drift in the direction of higher gas pressure, i.e. up the pressure gradient \\citep[e.g.][]{Klahr1997}. Besides the inward drift reduction in pressure bumps, dust grains do not experience the potentially damaging high-velocity collisions due to relative radial and azimuthal drift anymore, thus reducing the fragmentation problem as well.  Pressure bumps may themselves be caused by planets that have formed by a yet unknown mechanism. A possible first planet embryo can be explained by the growth of a seed (cm-sized object)  by sticking mutual collisions, including mass transfer and bouncing effects as possible outcomes of the collisions \\citep{windmark2012}. Once a massive planet is formed and opens a gap, a  pressure bump is expected in the radial direction  at the outer edge of the gap due to the depletion in the gas surface density. This pressure bump becomes a ``particle trap'', a region where particles accumulate and grow. This trap is therefore a good planet incubator for the next generation of planets. Particle trapping creates a ring-like emission at mm-wavelengths, whose radial position and structure depend on the disk viscosity, mass, and location of the planet \\citep{pinilla12a}. In addition, the non-axisymmetric emission of transition disks revealed by sub-mm observations \\citep{brown09} and recently confirmed with ALMA and CARMA observations \\citep{casassus13, vandermarel2013, isella2013, fukagawa2013, perez2014} may also be a result of particle trapping that lead to strong asymmetric maps \\citep{birnstiel2013, ataiee2013,  lyra2013}.  Spatially resolved, multi-frequency observations will provide crucial information about the dust density distribution of transition disks. Radial variations in the dust opacity have been revealed for individual disks \\citep[e.g.][]{perez2012, trotta2013}; however, a statistical study cannot have been done yet. In this paper, we show that in the case of transition disks,  a change in the cavity radius affects the spatially integrated spectral index, a quantity already available for a significant number of transition disks.  We present  3mm-observations of four transition disks: LkH$\\alpha$~330, UX~Tau~A, LRLL~31, and LRLL~67  carried out with Plateau de Bure Interferometer (PdBI).  The observations and data reduction of our four targets are presented in Sect.~\\ref{observations}, together with the information on additional transition disks. Section~\\ref{results} presents the integrated spectral index of our sample with the data collected from the literature. In this section a comparison between the spectral indices of transition and regular protoplanetary disks is presented, as well as a possible correlation between the cavity radii and the spectral index.  In section~\\ref{section_discussion} we compare the observations to two different kinds of models: first of all, dust evolution models based on planet-related cavities, where a pressure bump is formed at the outer edge of the gap, which traps millimetre particles; second, models with a power law size distribution of particles with an artificial cut at different radii to imitate dust cavities. Finally, in Sect.~\\ref{conclusion} we give the main conclusions of this work.  \\begin{table*} \\label{table1} \\centering \\tabcolsep=0.10cm \\begin{tabular}{c||cccccccccccccc} \\\\ \\textbf{{\\tiny Name}} & {\\tiny Ra} & {\\tiny Dec} &  {\\tiny SpT} & {\\tiny $T_{\\mathrm{eff}}$}&{\\tiny d}&{\\tiny $L_\\star$}&{\\tiny $R_\\star$}&{\\tiny $M_\\star$}&{\\tiny $\\dot{M}$}&{\\tiny $R_{\\mathrm{cav}}$} & {\\tiny $F_{\\mathrm{0.8mm}}$(mJy)}&{\\tiny $F_{\\mathrm{3.0mm}} $(mJy)}& {\\tiny $F_{\\mathrm{1.0mm}} $}&$\\alpha_{0.88-3\\mathrm{mm}}$\\\\ &{\\tiny [J2000]}&{\\tiny [J2000]}&&{\\tiny (K)}&{\\tiny (pc)}&{\\tiny ($L_\\odot$)}&{\\tiny ($R_\\odot$)}&{\\tiny ($M_\\odot$)}&{\\tiny ($M_\\odot~$yr$^{-1}$)}&{\\tiny (AU)}& {\\tiny$\\pm \\sigma$(mJy/beam)} & {\\tiny$\\pm \\sigma$(mJy/beam)} &{\\tiny (mJy)} &\\\\ \\hline \\hline \\\\ {\\tiny LkH$\\alpha$~330}&{\\tiny03 45 48.29}&{\\tiny +32 24 11.9} &{\\tiny G3}&{\\tiny5830}&{\\tiny 250}& {\\tiny 15}&{\\tiny 3.75}&{\\tiny 2.2}&{\\tiny 2$.0 \\times 10^{-9}$}&{\\tiny 68$^{R}$}& {\\tiny 210$\\pm$2.1}& {\\tiny 3.9$\\pm$0.12}& {\\tiny 138.6$\\pm$21.1}& {\\tiny 3.25$\\pm$0.40}\\\\ \\\\  {\\tiny UX~Tau~A}&{\\tiny 04 30 04.00}&{\\tiny +18 13 49.3} & {\\tiny G8}&{\\tiny5520}&{\\tiny 140}& {\\tiny 3.5}&{\\tiny 2.05}&{\\tiny 1.5}&{\\tiny $1.0\\times10^{-8}$}&{\\tiny 25$^{R}$}& {\\tiny 150$\\pm$1.5}& {\\tiny 10.1$\\pm$0.33}& {\\tiny 113.2$\\pm$17.2}& {\\tiny 2.20$\\pm$0.40} \\\\ \\\\  {\\tiny LRLL~31}&{\\tiny 03 44 18.00}&{\\tiny +32 04 57.0}&{\\tiny G6}&{\\tiny5700}&{\\tiny 315}& {\\tiny 5.0}&{\\tiny 2.3}&{\\tiny 1.6}&{\\tiny $1.4\\times10^{-8}$}&{\\tiny 14$^{S}$}& {\\tiny 62$\\pm$6.0}& {\\tiny 2.88$\\pm$0.08}& {\\tiny 45.0$\\pm$8.1}& {\\tiny 2.50$\\pm$0.44} \\\\ \\\\  {\\tiny LRLL~67} &{\\tiny 03 44 38.00} &{\\tiny +32 03 29.0} &{\\tiny M0.75}&{\\tiny3720}&{\\tiny 315}& {\\tiny 0.5}&{\\tiny 1.8}&{\\tiny 0.5}&{\\tiny $1.0\\times 10^{-10}$}&{\\tiny 10$^{S}$}& {\\tiny 25$\\pm$11}& {\\tiny 1.21$\\pm$0.09}& {\\tiny 18.2$\\pm$8.5}& {\\tiny 2.47$\\pm$0.92} \\\\ \\\\ \\hline \\end{tabular} \\caption{{\\bf Observed targets.} Column~1: Name of the disks observed with PdBI. Column~2 and~3: Right ascension and declination of the targets. Column~4: Spectral type. Column~5: Effective temperature of the central star. Column~6: Estimated distance to the target. Column~7, 8 and 9:  Stellar luminosity, radius and mass respectively. Column~10: Accretion rate. Column~11: Radial cavity extension ($R$ and $S$ implies that the cavity is either resolved or derived from SED modelling, respectively). Column~12: Flux at~345~GHz. The data is from \\cite{andrews11}  for LkH$\\alpha$~330 and UX~Tau~A, and from \\cite{espaillat12} for LRLL~31 and LRLL~67 and references therein. Column~13: Flux at~100~GHz from the PdBI observations. Column~14: Calculated flux at~300GHz ($\\sim~1$mm) and the error assuming an additional systematic uncertainty from  calibration of 15\\%. Column~15: Integrated spectral index calculated with the observed fluxes (columns 13 and 14).} \\end{table*} ",
        "conclusions": "\\label{conclusion}  We have reported 3mm observations carried out with PdBI of four transition disks LkH$\\alpha$~330, UX-Tau~A, LRLL~31, and LRLL~67.  Using previous observations of these targets at 880~$\\mu$m, we calculated the spectral index, and found that for  these sources the values for $\\alpha_{\\mathrm{mm}}$ are lower than typical values of the ISM dust. However, the disk-integrated spectral index $\\alpha_{\\mathrm{mm}}$ is close to the value for ISM-dust ($\\sim$~3.7) for LkH$\\alpha$~330 ($\\alpha_{\\mathrm{mm}}=3.25\\pm{0.40}$), which  is also the target with the largest inner hole \\citep[$\\sim$~68~AU,][]{andrews11b}. The sources UX-Tau~A and LRLL~31 have smaller cavity radii of 25~AU and 14~AU, respectively, and lower integrated spectral indices, $\\alpha_{\\mathrm{mm}}=2.20\\pm{0.40}$ for UX-Tau~A  and $\\alpha_{\\mathrm{mm}}=2.50\\pm{0.44}$ for LRLL~31. Owing to the large uncertainties on the millimetre flux of  LRLL~67, there is no conclusive correlation between the integrated spectral index and its cavity radius.  Recent observations at different wavelengths have shown that the opacity may  change radially throughout the disk \\citep{isella2010a, Guilloteau2011, perez2012, trotta2013}. Studies that combine dust evolution and planet-disk interaction for transition disks \\citep{pinilla12a} indicate  that the dust size distribution may present strong spatial variations. In this picture, the inner part is depleted of mm grains  leading to radial variations of the dust opacity. Thus,  the opacity index has higher values in the inner regions than  the outer regions, in contrast to observations of classical protoplanetary disks. From  models that combine hydrodynamical simulations and dust evolution models for the planet-related cavities, a linear relation between the disk integrated spectral  index and the cavity radius is inferred, because the millimetre emission is dominated by the dust at the outer edge of the cavity, i.e. in the pressure bump. This suggests that the spectral index is higher for larger cavities.  Since our own sample is too small to draw statistically significant conclusions, we compiled all the known millimetre fluxes from transition disk sources for which these data exist at least for two separated millimetre wavelengths. Together with our observations, we collecteded a total of twenty transition disks for which the cavity radii are known either by SED modelling or spatially resolved observations,  and we calculate the disk integrated spectral indices. We found that  the integrated spectral index is higher for transition disks than for regular protoplanetary disks.  Moreover, the probability of a positive trend between the spectral index and cavity size is $\\sim$95\\% (Fig.~\\ref{alpha_flux_Rcav2}). Two models are considered in this work to explain this positive correlation: first of all, coagulation and fragmentation of dust in a disk whose cavity is formed by a massive planet located at different positions; second, a disk with different truncation radii and with a power-law dust-size distribution where turbulence sets the maximum grain size. Both models can explain the trend, and multi-wavelength observations with high angular resolution  are needed to distinguish between the two cases.  Observations with new telescope arrays such as ALMA, EVLA, and NOEMA will provide  enough sensitivity and angular resolution to detect the small cavities of many transition disks and constrain  the dust opacity distribution of those disks in more detail.  Together with scattered imaging, these observations will help us understand the correlation between the spectral index and disk properties, such as the cavity size, hence, the environment where the planets are born."
    },
    "1402/1402.2644_arXiv.txt": {
        "abstract": "Very Long Baseline Interferometry (VLBI) allows for high-resolution and high-sensitivity observations of relativistic jets, that can reveal periodicities of several years in their structure. We perform an analysis of long-term VLBI data of the quasar S5 1928+738 in terms of a geometric model of a helical structure projected onto the plane of the sky. We monitor the direction of the jet axis through its inclination and position angles. We decompose the variation of the inclination of the inner $2$ milliarcseconds of the jet of S5 1928+738 into a periodic term with amplitude of $\\sim 0.89^{\\circ }$ and a linear decreasing trend with rate of $\\sim 0.05^{\\circ }\\text{yr}^{-1}$. We also decompose the variation of the position angle into a periodic term with amplitude of $\\sim 3.39^{\\circ }$ and a linear increasing trend with rate of $\\sim 0.24^{\\circ }\\text{yr}^{-1}$. We interpret the periodic components as arising from the orbital motion of a binary black hole inspiraling at the jet base and derive corrected values of the mass ratio and separation from the accumulated 18 years of VLBI data. Then we identify the linear trends in the variations as due to the slow reorientation of the spin of the jet emitter black hole induced by the spin-orbit precession and we determine the precession period $T_{\\mathrm{SO}}=4852 \\pm 646$ yr of the more massive black hole, acting as the jet emitter. Our study provides indications, for the first time from VLBI jet kinematics, for the spinning nature of the jet-emitting black hole. ",
        "introduction": "The existence of Supermassive Black Holes (SMBH) at the centre of galaxies, combined with the important role of galaxy mergers in the hierarchical galaxy formation models and the long time-scale of the SMBH merging process suggest that SMBHs often exist in pairs \\citep[][]{Blandford1986iau,SearleZinn1978,Komossa2006}. The binary passes through three evolutionary stages during the galaxy merger \\citep{Begelman1980}. At first the stellar halo of each galaxy interacts with the central black hole of the other galaxy via dynamical friction \\citep[see e.g.][]{BinneyTremaine1987}. Gradually the black holes sink into a common gravitational potential well, and a bound, widely separated SMBH binary forms \\citep[a detailed review of massive black hole binary evolution can be found in][]{MerrittMilos2005}. After this, a transition stage follows. The separation then decreases due to dynamical friction and beyond a transition radius the dissipation of energy and angular momentum becomes dominated by gravitational radiation. General relativistic effects during this inspiral stage are described in a post-Newtonian (PN) framework by expanding the field equations in terms of the PN parameter $\\varepsilon =Gmc^{-2}r^{-1}$ (where $m$ is the total mass, $r$ is the separation, $G$ is the gravitational constant, and $c$ is the speed of light). PN techniques are involved in the description of merging binary-dynamics if $0.001<\\varepsilon<0.1$, while in the plunge and the ring-down stages numerical calculations are required to characterize the final SMBH. The spins of the black holes generate corrections to the dynamics at $1.5$PN orders through the spin-orbit coupling, at $2$PN through the spin-spin coupling and at even higher orders through various post-Newtonian corrections to these couplings. The complicated PN dynamics in the presence of the spins is described in detail in e.g. \\citet{Barker1975,Barker1979}, \\citet{Kidder1995}, \\citet{Gergely2010b,Gergely2010a}. When the spin and orbital momentum vectors are not aligned, a spin-orbit type precession will cause the spin vectors to slowly rotate about the orbital angular momentum vector.  SMBHs are believed to be the main engine of the activity in Active Galactic Nuclei (AGN) \\citep[for a review see e.g.][]{Begelman1984}. Apart from efforts of identifying binary black hole candidates in kpc-scale separated binaries with spatially resolved AGN \\citep[e.g.][]{Komossa2003}, indirect approaches have been also employed for the cases when gravitational radiation already dominates the merger of a sub-pc-scale separated binary. It has been estimated from the jet power, that at least one merger with orbital period of the order of one year should be detected in future blind surveys (at a nominal sky coverage of $ 10^{4}$ deg$^{2}$ with $0.5$ mJy sensitivity) through the electromagnetic (EM) signature of the SMBH binary on the jet \\citep{Shaughnessy2011}, manifesting itself as an increase in the jet power when the binary inspirals toward its barycentre.  Relativistic bulk motion in the jet of AGN causes Doppler boosting of the synchrotron radiation of the relativistic particles. Very Long Baseline Interferometry (VLBI) observations reveal superluminally outward moving objects in a number of relativistic jets. Such superluminal motions are only apparent, indicating projection effects. These sources often show extreme variability throughout the EM spectrum. The radio variability can be explained in terms of flaring events following the ejection of jet-components or of a wiggling jet, implying a variable Doppler factor. AGN variability is however mostly non-periodic, hence detecting periodic behaviour of the radio light curve and precessing jets would strongly indicate a SMBH binary at the jet base \\citep[e.g. OJ287,][]{Villata1998}. While a plausible explanation for the precession of the jet is the existence of a merging SMBH binary at its base, alternative models allowing for a single black hole and external causes for the jet precession have been proposed \\citep[e.g. see the Discussion in][and references therein]{Britzen2010}.  The precession of a jet originating in a SMBH binary system was discussed in e.g. \\citet[][]{Roos1993,Romero2000,Britzen2001,Lobanov2005,Valtonen2012,Caproni2012}. Several models have been developed to investigate the interaction between a SMBH binary and its astrophysical environment, revealing the presence of the binary in the AGN \\citep[e.g.][]{MacFadyen2008,Tanaka2013}. Recently, the helical modulation of the large-scale radio jet of J1502+1115 was attributed to the effect of the tight pair in a triple supermassive black hole system \\citep{Deane2014}. Observational data have already been employed to identify precessing jet candidates. AGN observables allowed placing constraints on some of the SMBH binary parameters, like the total mass, the orbital period and the separation of the binary, e.g. Mrk 501, \\citet{Villata1999}; 3C 273, \\citet{Romero2000}; BL Lac,  \\citet{Stirling2003}; 3C 120,  \\citet{Caproni2004}; 3C 345, \\citet{Lobanov2005}; S5 1803+784, \\citet{Roland2008}; NGC 4151, \\citet{Bon2012}. The effect of the spin was not included in the calculations of these works, and the binary motion was approximated by pure Keplerian orbits.  The extreme Kerr-limit of the rotating SMBH is characterized by the value $1$ of the dimensionless spin parameter $\\chi$ (the spin $S$ can be calculated as $S= G c^{-1} m^2 \\chi$). However, as the accretion disk radiates and some of that radiation is trapped by the black hole, the extreme limit is actually not reached and rather the so-called canonical spin limit of $\\chi _{\\mathrm{can}}\\simeq 0.998$ applies \\citep{Thorne1974}. A closed magnetic field line topology further connects the black hole horizon to the accretion disk, field lines acting as anchor chains, which further reduces the spin to $\\chi \\simeq 0.89$ \\citep{Kovacs2011}. Nevertheless this value is still extremely high, and as the presence of the jet signals a rich black hole environment, supporting the assumption of accretion, it is plausible to assume that at least the jet emitter black hole spins fast, and as such, the spin effects in the binary dynamics are non-negligible. The spin-powered jets are believed to be aligned with the black hole spin axis \\citep{Blandford1977}, therefore precessing jets can reveal the presence of the spin.  In this paper we investigate the source S5 1928+738. This is a core-dominated quasar  at a redshift $z=0.302$, its luminosity distance is $D_{\\mathrm{L}}=1620$ Mpc, and at this distance the spatial resolution is  $4.6$ pc/milliarcsecond \\citep[where the cosmological parameters are $\\Omega _{\\mathrm{m}}=0.314$, $\\Omega _{\\lambda}=0.686$, $H_{0}=67.4$ km s$^{-1}$ Mpc$^{-1}$,][]{Planck2013}. The source is a well studied member of a complete sample of the extragalactic flat spectrum radio sources taken from the S5 strong source survey \\citep{Kuhr1981}. The VLA map at a wavelength of $20$ cm \\citep{Johnston1987} reveals a structure comprising two lobes and a bright radio-core. Zooming into the jet base a one-sided core-jet structure appears on the milliarcsecond (mas) scale. The maximal apparent velocity detected in the jet, $8.1c$ \\citep[][]{Lister2013} suggests high-speed bulk motion that is nearly head on, the mean jet direction is pointing close to the observer line of sight (LOS). This feature makes the jet a strongly Doppler-boosted one.  An analysis of the arcsecond- and the mas-scale jet structure has been published by \\citet{Hummel1992} and the jet structure has been found to exhibit wiggles on both scales. A moving sine wave was fitted to the temporal motion of the components in the jet. \\citet{Roos1993} proposed that the host galaxy of S5 1928+738 harbours a binary SMBH at its centre, inducing the periodic structures detected in the jet. A total mass of $10^{8}$ M$_{\\odot} $ was adopted, which would account for the bolometric luminosity on the basis of the Eddington limit. The orbital period was calculated to be $\\sim 2.9$ years, the orbital separation $\\sim 0.005$ pc and the mass ratio $>0.1$. It has also been noted that if the wiggles in the mas-scale jet of S5 1928+738 are due to the orbital motion of a massive binary, the mean direction of the jet should obey a Lense-Thirring precession \\citep{LenseThirring1918} about the orbital angular momentum of the binary, with a period of the order of $10^{3}$ yr. \\citet{Murphy2009} used the VLBI Space Observatory Programme data to further constrain the binary parameters by fitting their equations to jet component data from $8$ epochs of observation performed at 5 GHz. The derived parameters are however affected by their assumption of the jet being emitted by the fictitious reduced mass particle.  The data collected on S5 1928+738 within the framework of the Monitoring Of Jets in Active galactic nuclei with VLBA Experiments (MOJAVE) survey span almost twenty years \\citep{Lister2009}. Careful analysis of this dataset allows us to detect and quantitatively analyse for the first time the spin-orbit precession of the jet emitter SMBH, manifesting itself as a small change over time in the direction of the jet axis. The spatial-length of the jet based on the MOJAVE-VLBI maps at $15$ GHz and assuming an inclination of $7\\degr$ is approximately $560$ pc.  The paper is organised as follows. In Section 2 we summarize the observational techniques which can reveal periodicities in the extragalactic jets observed at radio wavelengths. More specifically we present methods based on the investigation of the AGN jet morphology, AGN jet kinematics and AGN jet flux density variability. In Section 3 we present the analysis of the archival MOJAVE data of S5 1928+738, investigating the temporal evolution of the jet structure. In Section 4, we describe a geometrical model for the three-dimensional (3D) helical jet structure. By projecting the spatial jet onto the plane of the sky and employing the maximal apparent velocity seen in the jet (which provides information about the inclination angle), we constrain the model by fitting it to the MOJAVE observations on the jet. Based on the jet modelling and on the independent total mass estimate, we give the binary parameters in Section 5. In Section 6 we discuss our findings and give our conclusions. ",
        "conclusions": "In this paper we investigated the radio jet of S5 1928+738 based on calibrated data of the MOJAVE survey with focus on the perturbed jet ejection. Our analysis basically confirmed the model of \\citet{Roos1993} which in turn was based on the jet analysis by \\citet{Hummel1992} of a hidden SMBH binary at the jet base, causing the observed wiggling of the jet. The improvement upon this model, advanced in the present paper is the inclusion of the spin of the jet-producing black hole.  We developed a detailed geometrical model to describe the inner $2$ mas of the jet by using the $43$ GHz VLBA-map of \\citet{Lister2000}. From that we obtained the inclination and position angle of the jet axis and the intrinsic half-opening angle of the conical helix. Then we fitted the geometric model to the $15$ GHz VLBA data of almost twenty years. Three features of the jet have been revealed; (i) a conical helix shape, (ii) a periodical change in the direction of its symmetry axis, (iii) a slow additional reorientation in the average direction of the jet. Our spinning binary black hole model naturally explained the simultaneous presence of properties (ii) and (iii), and (i) was attributed to the presence of helical kink instabilities.  Jet precession (iii) can be induced through the Bardeen-Petterson interaction between a viscous accretion disk and a spinning black hole \\citep{BardeenPetterson1975}, if the disk is misaligned compared to the equatorial plane of the black hole. % However such a scenario cannot explain both (ii) and (iii). In an alternative scenario, the influence of the immediate environment could cause the wiggling of the jet. However, to explain the observed periodicities would require properly fine-tuned structures. So far there is no evidence for such periodic distribution of dense material close to the jet.  We adopted the total mass of the binary $m=8.13 \\times 10^{8}M_{\\odot }$  determined by \\citet{Woo2002} from the black hole mass-AGN continuum luminosity scaling relation (this value is different by a factor of $8.13$ from the one adopted by \\citealt{Roos1993}). The helical jet model with periodic jet axis and VLBI data implied the orbital period $T=4.78\\pm 0.14$ yr (this is a factor of $1.59$  larger than the value given by \\citealt{Roos1993}). With these values we calculated the binary separation as $r=0.0128\\pm 0.0003$ pc. These parameters imply that the SMBH binary is in the inspiral phase, but far from coalescence, with post-Newtonian parameter $\\varepsilon \\approx 0.003$.  Long-term monitoring of the radio jet allowed us to identify the linear trend in the evolution of the inclination and position angles of the jet axis, interpreted as arising from the spin-orbit precession of the jet emitter SMBH. The mass ratio most likely falls into the range $\\nu \\in \\left[ 0.21:1/3 \\right] $. The spin-orbit precession period was identified as $T_{\\mathrm{SO}}=4852 \\pm 646$ yr and the gravitational lifetime emerged as $T_{\\mathrm{merger}}=(1.44 \\pm 0.19)\\times 10^6$ yr.  Although we cannot rule out that other models could explain the observed jet structure, we showed that the VLBI data of S5 1928+738, extending over almost twenty years is consistent with the model of a spinning binary black hole lying at the jet base, where the larger black hole has a spin detectable through its spin-orbit precession. Measurements of slow increase in the average flux density of the jet in the newest three epochs (2013.34, 2013.58 and 2013.96) further support the model, as such an increase is predicted by spin-orbit precession. Our study thus provides indications, for the first time from VLBI jet kinematics, for the spinning nature of the jet-emitting black hole.  As the \\-MOJAVE survey is still ongoing, further data on S5 1928+738 may better constrain the parameters of the model. With a significantly increased amount of data in principle it would be possible to monitor higher order post-Newtonian effects too. Beyond spin induced precession, such an analysis could also reveal the magnitude of the spin, unavailable at the accuracy of the present analysis."
    },
    "1402/1402.2702_arXiv.txt": {
        "abstract": "We present a detailed comparison between the photometric properties of the bulges of two simulated galaxies and those of a uniform sample of observed galaxies. This analysis shows that the simulated galaxies have bulges with realistic surface brightnesses for their sizes and magnitude. These two field disc galaxies have rotational velocities $\\sim$ 100 km/s and were integrated to a redshift of zero in a fully cosmological $\\Lambda$ cold dark matter context as part of high-resolution smoothed particle hydrodynamic simulations. We performed bulge-disc decompositions of the galaxies using artificial observations, in order to conduct a fair comparison to observations. We also dynamically decomposed the galaxies and compared the star formation histories of the bulges to those of the entire galaxies. These star formation histories showed that the bulges were primarily formed before {\\em z} = 1 and during periods of rapid star formation. Both galaxies have large amounts of early star formation, which is likely related to the relatively high bulge-to-disc ratios also measured for them. Unlike almost all previous cosmological simulations, the realistically concentrated bulges of these galaxies do not lead to unphysically high rotational velocities, causing them to naturally lie along the observed Tully--Fisher relation. ",
        "introduction": "Computational astronomers have made great strides in reproducing the observed disc structure of spiral galaxies in $\\Lambda$ cold dark matter simulations, as studies of the Tully--Fisher relation \\citep{Robertson04, Governato07, Stinson10, Piontek11}, angular-momentum content \\citep{Scannapieco09}, and disc sizes \\citep{Brooks11} have shown. Historically, though, the corresponding bulges of simulated galaxies were much too large and concentrated compared to observed bulges \\citep{Bullock01, vandenBosch01b, Binney01, vdb02, donghia06, dn07,Stinson10, scannapieco11}. This difficulty in reproducing the bulges of disc galaxies represents the more-persistent aspect of the `angular-momentum catastrophe' \\citep{Navarro94}.  In the past several years, only a handful of simulations have been able to combine the resolution necessary to resolve bulges in cosmological simulations with star formation and feedback models capable of producing galaxies with realistic mass bulges. These simulations have variously succeeded by limiting the star formation efficiency \\citep{Agertz10}, pre-heating the gas through early stellar feedback \\citep{Stinson13}, or increasing the strength of supernova-driven (SN-driven) outflows, either by scaling a kinetic wind model \\citep{Okamoto12, Vogelsberger13} or concentrating the energy through a high star formation threshold \\citep{Guedes11, Brook11a}. Similar strong outflows have been shown to reduce the amount of low-angular momentum material \\citep{Pontzen11}, resulting in bulgeless or cored dwarf galaxies \\citep{Oh11, Governato12,Teyssier13} and disc galaxies with realistic, rising rotation curves \\citep{Christensen12a, AnglesAlcazar13}. % Despite these improvements, no single method has been shown to reproduce the observed properties of bulges in a range of galaxies. As such, the masses and concentration of bulges remain as firm limits on the formation of galaxies. In particular, the observed scaling relations between bulge surface density, magnitude, and size encapsulate much of the observational constraints.  Given the importance of reproducing bulge properties, it is critical to analyse them in an observationally motivated fashion. Here we undertake a careful photometric analysis of the bulges of two simulated galaxies and compare them to an observed sample analysed in an identical fashion. These galaxies were previously shown to have flat rotation curves and a reduced mass of low-angular momentum material as the result of feedback-driven outflows \\citep[][hereafter C12a]{Christensen12a}. In this Letter, we show that they also have realistically concentrated bulges. ",
        "conclusions": "\\label{sec:res5} In this Letter, we have taken an observationally motivated approach to analyzing the bulges of two simulated galaxies, previously shown to have rising rotation curves (C12a). We photometrically decomposed H-band images into bulge and disc components and compared the bulge properties to an observational sample of disc and elliptical galaxies. We determined that these simulations had bulges of the appropriate surface brightness for their magnitude and size. This success indicates that the centres of the simulated galaxies have appropriate stellar distributions and that they are not overly concentrated. Their relatively high bulge-to-total ratios, however, remain a concern and demonstrate a need for a further reduction in the central stellar mass. In particular, reducing the amount of early star formation could both lower the bulge mass and result in more realistic SFHs.  We compared the SFHs of the kinematically selected bulges to those of the entire galaxies. We found that the bulges formed primarily during the first half of the galaxies' lifetimes and that bulge star formation occurred during periods of peak star formation, possibly driven by mergers. Of the two galaxies, the SFHs and photometric properties of one were most consistent with the properties of classical bulges, while the other shared characteristics with both classical and pseudo-bulges. Finally, we verified the global structure of the simulated galaxies by comparing them to the observed baryonic and stellar Tully--Fisher relations."
    },
    "1402/1402.1910_arXiv.txt": {
        "abstract": "Given an approximately centered image of a spiral galaxy, we describe an entirely automated method that finds, centers, and sizes the galaxy and then automatically extracts structural information about the spiral arms. For each arm segment found, we list the pixels in that segment and perform a least-squares fit of a logarithmic spiral arc to the pixels in the segment. The algorithm takes about 1 minute per galaxy, and can easily be scaled using parallelism. We have run it on all $\\sim$644,000 Sloan objects classified as ``galaxy'' and large enough to observe some structure. Our algorithm is stable in the sense that the statistics across a large sample of galaxies vary smoothly based on algorithmic parameters, although results for individual galaxies can sometimes vary in a non-smooth but easily understood manner. We find a very good correlation between our quantitative description of spiral structure and the qualitative description provided by humans via Galaxy Zoo. In addition, we find that pitch angle often varies significantly segment-to-segment in a single spiral galaxy, making it difficult to define ``the'' pitch angle for a single galaxy. Finally, we point out how complex arm structure (even of long arms) can lead to ambiguity in defining what an ``arm'' is, leading us to prefer the term ``arm segments''. ",
        "introduction": "\\label{sec:Intro} \\vspace{-3mm}  The Hubble Ultra Deep Field represents about 1/13,000,000 of the celestial sphere and contains about 10,000 galaxies, suggesting the entire sky contains upwards of $10^{11}$ galaxies. Gaining quantitative structural information for this number of galaxies will require automated methods. For spiral galaxies, existing methods for visual classification \\cite{Hubble1936,DeVaucouleurs1959} are either subjective or non-quantitative, while currently available semi-automated methods \\cite{DeSouza2004,Simard1998,Peng2002a,Peng2010,BenDavis2012,Ma2001,Au2006,Ripley1990,Shamir2011}. are either too simplistic or require significant human input.  \\begin{figure*}[t] \\plotsix{2-autoCrop.eps}{3-contrastEnhanced.eps}{4-orientationField_32x32-from-64x64.eps}{5-clustersAndArcs_arcEnhanced.eps}{6-mergedClustersAndArcs_arcEnhanced.eps}{7-arcOverlayOnAutoCrop_arcEnhanced.eps} $\\begin{array}{cccccc} \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\,  ({\\bf a}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; &\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\, ({\\bf b}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; &\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; ({\\bf c}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; &\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; ({\\bf d}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; &\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; ({\\bf e}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; &\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; ({\\bf f}) \\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\;\\; \\end{array}$ \\caption{Steps in describing a spiral galaxy image. {\\bf a)} The centered and de-projected image. {\\bf b)} Contrast-enhanced image. {\\bf c)} Orientation field (at reduced resolution for pedagogical reasons). {\\bf d)} Initial arm segments found via Hierarchical Agglomerative Clustering (HAC) of nearby pixels with similar orientations and consistent logarithmic spiral shape, and the associated logarithmic spiral arcs fitted to these clusters. {\\bf e)} Final pixel clusters (and associated arcs) found by merging compatible arcs. {\\bf f)} Final arcs superimposed on image (a).  Red arcs wind S-wise, cyan Z-wise. } \\label{fig:parse} \\end{figure*} ",
        "conclusions": "\\vspace{-3mm}  We have run our code on every item in the Sloan Digital Sky Survey (SDSS) that is classified as ``galaxy''.  Unfortunately, SDSS does not distinguish between spiral and non-spiral galaxies.  We are currently working on a machine learning algorithm that uses the output of our code to distinguish between spiral and non-spiral galaxies.  Preliminary results are encouraging (agreeing approximately 90--95\\% with Galaxy Zoo humans), and will be presented in a future paper.  Once we can separate out spiral galaxies, further studies will be performed concerning how spiral structure correlates with other variables such as color, redshift, local environment, etc.  Code and data are available from WBH upon request.  \\vspace*{-3mm}"
    },
    "1402/1402.3915_arXiv.txt": {
        "abstract": "In observations of flare-heated electrons in the solar corona, a longstanding problem is the unexplained prolonged lifetime of the electrons compared to their transit time across the source. This suggests confinement. Recent particle-in-cell (PIC) simulations, which explored the transport of pre-accelerated hot electrons through ambient cold plasma, showed that the formation of a highly localized electrostatic potential drop, in the form of a double layer (DL), significantly inhibited the transport of hot electrons (T.C. Li, J.F. Drake, and M. Swisdak, 2012, ApJ, 757, 20). The effectiveness of confinement by a DL is linked to the strength of the DL as defined by its potential drop. In this work, we investigate the scaling of the DL strength with the hot electron temperature by PIC simulations, and find a linear scaling. We demonstrate that the strength is limited by the formation of parallel shocks. Based on this, we analytically determine the maximum DL strength, and find also a linear scaling with the hot electron temperature. The DL strength obtained from the analytic calculation is comparable to that from the simulations. At the maximum strength, the DL is capable of confining a significant fraction of hot electrons in the source. ",
        "introduction": "INTRODUCTION}     Coronal electrons can be accelerated to very high energies, over 100 keV \\citep{Lin03, Krucker10, Krucker07,Tomczak09}, during flares and coronal mass ejections. Acceleration is believed to occur high in the corona \\citep{Fletcher11}. Hard and soft X-rays emitted by the accelerated electrons are observed in both the corona and chromosphere. Emission from the chromosphere is usually brighter because of the higher density, which allows for more efficient interaction with the X-ray-producing energetic electrons. Improved detection sensitivity, however, has led to more measurements of X-ray sources in the corona. These sources are situated at or distinctly above the top of the soft X-ray flaring loops. A comprehensive analysis suggests that above-the-looptop sources are a common feature of all flares \\citep{Petrosian02}.   The transport of energetic electrons from the corona to the chromosphere is crucial to understanding energy release in flares. Transport effects can modify the energy distribution of the propagating electrons and hence the observed X-ray spectra, affecting the interpretation of acceleration models. There is evidence for both free-propagation and confinement of energetic electrons. Evidence for free-propagation of energetic electrons, i.e., no interaction with the ambient plasma, was reported in previous time-of-flight measurements of hard X-ray emission \\citep{Asch95,Aschwanden96,AschSchwartz95}. A systematic time delay between lower-energy hard X-rays with respect to the higher energies was measured from a large sample (over 600) of solar flares. The time-of-flight delays indicate lower-energy electrons arrive at the chromosphere after those with higher energy as they freely stream down the flare loop from the corona. On the other hand, observations of above-the-looptop sources reveal that the lifetime of the energetic electrons is two orders of magnitude longer than their free-streaming transit time across the source \\citep{Masuda94, Krucker10, Krucker07}. This requires confinement of the electrons in the source region. Recently, a systematic study of solar flares with both looptop and footpoint emissions found that the number of electrons required to explain observations is 2-8 times higher at the looptop than at the footpoint \\citep{Simoes13}. This suggests electron accumulation at the looptop. Another systematic study reported that the difference in spectral index between non-thermal coronal and footpoint emission in some events is considerably greater than 2 - which is expected for the simple case that the same electron population, which produces looptop emission through thin target bremsstrahlung because of the lower density in the corona, free-streams to the footpoints to emit through thick target bremsstrahlung due to the higher density in the chromosphere \\citep{Battaglia06}. This requires a filter effect in the propagation preferentially reducing the distribution at lower energies. Such a filter can be an electric field. An important question on electron transport is why some electrons appear to freely propagate while some are filtered or even trapped. The physics of this subject, however, remains poorly understood.    The transport of electron heat flux has previously been modeled as anomalous conduction \\citep{Manheimer77, Tsytovich71} in which anomalous resistivity arises from electron scattering by turbulent wave fields. Turbulence is excited by instabilities that result from the interaction between energetic electrons and ambient plasma. In a similar model, it was suggested that ion-acoustic turbulence driven by a return current in response to hot electron streaming and an induced polarization electric field would drive a potential that would suppress electron transport \\citep{Levin93}. An anomalous conduction front that moved along the flare loop at the head of an expanding hot electron source was considered as a means to confine hot electrons for the production of hard X-rays \\citep{Smith79}. The model, based on a one-dimensional (1D) one-fluid code, however, did not resolve ion inertial lengths, let alone capture processes occurring at electron scales. Later, 1D electrostatic particle-in-cell (PIC) simulation that resolved the shortest electron scale, the Debye length, did not reveal a conduction front \\citep{McKean90}. Recently, the existence of a thermal front, literally defined as a region that links plasmas in thermal nonequilibrium and sustains the temperature difference for longer than the electron free-streaming time, was studied in 1D electrostatic Vlasov simulations \\citep{Arber09}. The formation of a temperature difference propagating at a speed comparable to the ion acoustic speed was observed and the behavior was identified as a conduction front. The physics of the responsible mechanism was not identified or investigated. We suggest, on the basis of the simulations and analysis in this paper, that the propagating front is an ion acoustic shock and the temperature difference to be a result of shock heating due to its extremely sharp transition. Other observed non-propagating temperature jumps that were associated with a potential jump were likely DLs, as mentioned by the authors.    More recently, the transport of coronal energetic electrons was studied by electromagnetic 2D PIC simulations \\citep{Li12} (hereafter as LDS). It was shown that transport suppression began as the electrons propagated away from the acceleration site. The suppression was caused by the formation of a DL and associated potential barrier that reflected electrons back to the source region. The DL reduced the electron heat flux by nearly 50\\%.     A DL is a localized region that sustains a potential drop in collisionless plasmas \\citep{Block78,Raadu88}. The potential drop comes from a large-amplitude electrostatic electric field sandwiched between two adjacent layers of equal and opposite charges. The structure is globally neutral, but quasi-neutrality is locally violated within it, which occurs at scales of $\\sim$10 Debye lengths $\\lambda_{De}$. An ideal DL is a monopolar electric field, but more generally, a DL can be bipolar. Therefore, instead of a monotonic drop in the potential $\\phi$, there can be dips or bumps at the low or high potential sides, but an overall potential drop $\\phi_{DL}$ across the structure. $\\phi_{DL}$ is the measure of the strength of a DL. A DL reflects particles if $\\phi_{DL}$ is greater than their kinetic energy. In LDS, it is the reflection of energetic electrons by a DL that reduces their transport from the source.   Numerical simulations with a kinetic model have been used to study DL dynamics since they can both resolve kinetic scales down to $\\lambda_{De}$, where DLs occur, and probe nonlinear phenomena. Methods used to generate DLs include a current and an applied potential. In LDS and this work, the DL forms as a result of strong currents driven by an imposed field-aligned temperature jump in the initial state. Earlier 1D PIC simulations of systems driven by a subthermal electron current showed that ion trapping in the potential dip of a (bipolar) DL was associated with decay of a DL \\citep{Barnes85}. The trapping of ions slowed down the movement of the DL structure as the trapped ions added an inertial drag to it \\citep{Chanteur83}. It was argued that such slowing down led to the decay of the DL \\citep{Barnes85}. On the other hand, 1D PIC simulations with an applied potential reported that a DL decayed as a result of the propagation of solitons, which developed from the DL structure, across the DL width \\citep{Sato81}. Only very limited work exists on the saturation of DLs.     The degree of transport suppression depends on the maximum strength of the DL at saturation. In LDS, the generation mechanism of a DL was studied and identified as the Buneman instability involving the ions and return current electrons. LDS showed that the potential increased with increasing hot electron temperature, but the functional dependence on this temperature was not studied. In this work, we investigate the saturation mechanism of the DL with simulations and analytic modeling. We carry out a series of simulations with different initial values of the hot electron temperature and show that the DL potential at saturation scales linearly with the hot electron temperature. The DL saturates when its potential jump is large enough to accelerate ions above the local sound speed. The result is a parallel ion acoustic shock that stabilizes the Buneman instability and therefore saturates the DL. We demonstrate that the shock formation criterion predicts a maximum DL strength that is proportional to the hot electron temperature in agreement with simulations. At its maximum strength, the DL is observed to reflect, and hence contain, a significant fraction of electrons in the source.     In addition, we observe that an anisotropy with $T_\\perp > T_\\parallel$ develops in the hot electrons that pass through a DL. Such an anisotropic (pancake) distribution favors further hot electron trapping in the presence of a magnetic mirror. Since the magnetic geometry of flaring loops resembles a mirror configuration, we explore combining a magnetic mirror with the DL. We demonstrate that the combination significantly enhances confinement.    An outline of the work is as follows: the setup of the simulations and parameters used for flare settings are described in Section \\ref{sim}; in Section \\ref{result}, results from the simulations and the shock model are presented, and we provide evidence for the model and determine the saturation amplitude of a DL; in Section \\ref{mu_anis}, we combine the DL with a magnetic mirror and show that further suppression of transport is obtained; further discussion of our results is given in Section \\ref{dis}; and we summarize in Section \\ref{con}. ",
        "conclusions": ""
    },
    "1402/1402.1584_arXiv.txt": {
        "abstract": "{Gravity modifies the spectral features of young brown dwarfs (BDs).  A proper characterization of these objects is crucial for the  identification of the least massive, and latest-type objects in star-forming regions, and to explain the origin(s) of the peculiar spectro-photometric properties of young directly imaged extrasolar planets and BD companions.} {We obtained medium-resolution (R$\\sim$1500-1700) near-infrared (1.1-2.5 $\\mu$m) spectra of seven  young M9.5-L3 dwarfs classified at optical wavelengths. We aim to empirically confirm   the low surface gravity of the objects in the near-infrared. We also  test whether self-consistent atmospheric models correctly represent the formation and the settling of dust clouds in the atmosphere of young late-M and L dwarfs.} {We  used ISAAC (Infrared Spectrometer And Array Camera) at VLT (Very Large Telescope) to obtain the spectra of the targets. We compared them  to those of mature and young BDs, and young late-type companions to nearby stars with known ages, in order to identify and study gravity-sensitive features. We computed  spectral indices weakly sensitive to the surface gravity to derive near-infrared  spectral types. Finally, we found the best fit between each spectrum and   synthetic spectra from the  BT-Settl 2010 and 2013 atmospheric models. Using the best fit, we derived the atmospheric parameters of the objects and identify which spectral characteristics the models do not reproduce.} {We confirmed that our objects are young BDs and we found near-infrared spectral types in agreement with the ones determined at optical wavelengths.  The spectrum of the $\\mathrm{L2{\\gamma}}$ dwarf \\object{2MASSJ232252.99-615127.5} reproduces well the spectrum of the planetary mass companion \\object{1RXS J160929.1-210524b}.  BT-Settl models fit the spectra and the 1-5 $\\mu$m spectral energy distribution of the L0-L3 dwarfs for temperatures between 1600-2000 K. But the models fail to reproduce the shape of the H~band, and the near-infrared slope of some of our targets. This fact, and the best fit solutions found with super-solar metallicity are indicative of a lack of dust, in particular at high altitude,  in the cloud models.} {The modeling of the vertical mixing and of the grain growth will be revised in the next version of the BT-Settl models. These revisions may suppress the remaining non-reproducibilities.  Our spectra provide additional templates for the characterization of the numerous young L-type companions that will be detected in the coming years by planet imaging instruments such as VLT/SPHERE, Gemini/GPI, Subaru/SCexAO, and LBTI/LMIRCam.} ",
        "introduction": "Since the discovery of the first substellar objects (\\citealt{1995Natur.378..463N,Rebolo1995}), large infrared (IR) surveys have unearthed hundreds of brown dwarfs (BDs) in the field and in star forming regions. The spectral energy distributions of BDs peak in the infrared. Their spectra are dominated by broad, overlapping condensate and molecular absorption features \\citep[see][]{Kirkpatrick}.  The strength of these features depends on a combination of photospheric temperature, gas, pressure and dust properties, which in turn are related to the effective temperature ($\\mathrm{T_{eff}}$), surface gravity ($g$), and metallicity of the objects. Spectroscopic studies of mature ($\\gg 400$ Myr) field BDs led to the definition of the spectral classes `L' \\citep{1999ApJ...519..802K, 1999AJ....118.2466M} for objects dominated by H$_{2}$O, FeH, and CO absorption bands in the near-infrared (NIR), and `T'  \\citep{2002ApJ...564..421B, 2002ApJ...564..466G} for objects which exhibit strong CH$_{4}$+H$_{2}$O bands and collision-induced absorption (CIA) due to H$_{2}$ at these wavelengths. The change in the spectral morphology,  used for the classification, is mostly tied to changes in $\\mathrm{T_{eff}}$ \\citep[e.g.,][]{2009A&A...503..639T}.  Substellar objects contract and cool down with time \\citep{1997ApJ...491..856B, BCAH98}. Therefore, young objects can have larger radii and higher luminosities than older and more massive objects despite an identical effective temperature. The lower surface gravity of young objects ($g \\varpropto M/R^{2}$) can be directly accessed by  observation and be used to break the degeneracy.  Low surface gravity results in peculiar  spectral characteristics such as the triangular H band shape in the near-infrared \\citep{Zapatero-Osorio2000, Lucas}, and reduced alkali lines in the optical and near-infrared \\citep[e.g.,][]{McGovern, Cruz}.  These modifications have nonetheless mostly been investigated for objects earlier than $\\sim$L4 found in young clusters and star forming regions \\citep[e.g.,][]{Luhman1997, 1999ApJ...525..466L, Lucas,  2004ApJ...617..565L, 2004AJ....127..449M, Luhman2004, McGovern, Lodieu, Weights, 2012A&A...539A.151A}.  Advances in  high contrast and high resolution imaging in the near-infrared (1-5 $\\mu$m; NIR) led to the discovery of late-type companions to young nearby stars  straddling the planet/BD boundary \\citep[e.g.,][]{2004A&A...425L..29C, 2005A&A...438L..29C, 2008ApJ...689L.153L, 2008Sci...322.1348M, 2010Natur.468.1080M, 2010Sci...329...57L, 2011ApJ...726..113I, 2013A&A...553L...5D, 2013ApJ...772L..15R, 2013ApJ...774...11K, 2013arXiv1310.4825C}. Most of these objects have estimated effective temperatures similar to those of L dwarfs. The low to medium-resolution (R$\\leq$5000) spectra and infrared photometry of these objects exhibit peculiar features \\citep[red pseudo-continua, triangular H~band shape, lack of methane absorption, reduced K I and Na I lines; ][]{2008ApJ...689L.153L, Bonnefoy2010, Patience, 2011ApJ...729..139W, Barman, Bonnefoy2013, 2013ApJ...768...24O}. These peculiarities are likely  directly related to the expected low surface gravity of the objects.  Nevertheless, the  currently limited sample of young late-type objects with high S/N (signal-to-noise) spectra make the establishment of a proper empirical classification scheme challenging for these objects. \\citet{Cruz} have identified a population of young bright and nearby  L dwarfs isolated in the field. They  developed a classification scheme in the optical. This scheme has been extended in the NIR for young L dwarfs by \\citet{Bonnefoy2013} and \\citet{Allers2013}. Several of these peculiar L-type dwarfs share similarities with the spectra of young companions at these wavelengths \\citep{Bonnefoy2010, Bonnefoy2013, Allers2013, 2013AJ....145....2F, 2013arXiv1310.0457L}.  Getting more NIR spectra of young isolated objects are therefore needed to consolidate the classification scheme and to identify  further analogues to directly imaged exoplanets.  Atmospheric models  allow us to disentangle the effect of varying $T_{eff}$, log~\\textit{g}, and (metallicity) \\textit{M/H} on the spectral features. Below $T_{eff}$ $\\sim$~2600~K, models predict that clouds of iron and silicate grains begin to form, changing the opacity (\\citealt{Lunine}, \\citealt{Tsuji}, \\citealt{Burrows_Sharp}, \\citealt{Lodders}, \\citealt{Marley2000}, \\citealt{Marley2001}, \\citealt{Allard2001}). The formation and the gravitational sedimentation of these dust clouds are  influenced by the surface gravity. Dust cloud formation is  expected to be more efficient at low gravity because the atmosphere is more extended, and the gas cooler. Low gravity tends to make the convection and the resulting mixing more efficient as well.  Modeling of the spectro-photometric data on young L and early-T type companions with parametrized models (\\citealt{Marley_Ackerman}, \\citealt{Burrows}) has revealed anomalously thick clouds (\\citealt{Barman2011a},  \\citealt{Barman2011b}, \\citealt{Skemer}, \\citealt{Marley2013}). These peculiar cloud properties may explain why some companions are \"underluminous\" in some bands (\\citealt{Skemer}, \\citealt{Marley2012}). Self-consistent atmospheric models, such as the BT-Settl models \\citep{Allard} and the Drift-PHOENIX models, \\citep{Helling} use cloud models where the dust properties do not require defining any additional free parameters other than log \\textit{g}, $T_{eff}$, \\textit{M/H}. Synthetic spectra for a specific set of atmospheric parameters can be compared to empirical spectra. These models are just starting to be tested  on spectra of young late-type objects \\citep[companions and free-floating objects;][]{Bonnefoy2010, 2011A&A...529A..44W, 2012A&A...540A..85P, Bonnefoy2013}. We used them on spectra of young M5.5-L0 dwarfs in \\citet{Bonnefoy2013}  to reveal a drop of the effective temperature at the M/L transition.  We suggested that this drop could be induced by an improper handling of the formation of dust clouds at the M/L transition in the models. The test however could not be extended to later spectral types and lower effective temperatures due to the lack of a consistent sample of young objects in the L dwarf regime at that time.  In this paper, we present a homogeneous set of seven medium-resolution (R$\\sim$1500-1700) spectra of M9.5-L3 dwarfs, all classified at optical wavelengths. Our sample is composed of the M9.5 object \\object{DENIS-P J124514.1-442907} (also called TWA 29; hereafter DENIS J1245) a member of TW-Hydrae (5-10~Myr), and  the L0 dwarf  \\object{Cha J1305-7739} \\citep[][hereafter Cha 1305]{Jayawardhana}, one of the least massive objects of the Chameleon II cluster. We also present the spectra of five L dwarfs with features indicative of low surface gravity (L$\\gamma$ dwarfs) in the optical, identified by \\cite{Cruz}. These objects are the two L0$\\gamma$ dwarfs \\object{EROS~J0032-4405} \\citep[][hereafter EROS~J0032]{Goldman} and  \\object{2MASS J22134491-2136079} \\citep[][hereafter 2M2213]{Cruz}, the  L2$\\gamma$ dwarf \\object{2MASSJ232252.99-615127.5} \\citep[][hereafter 2M2322]{Cruz}, and the two  L3$\\gamma$ dwarfs \\object{2MASS J212650.40-814029.3} \\citep[][2M2126]{Cruz} and \\object{2MASSJ220813.63+292121.5} \\citep[][hereafter 2M2208]{Cruz}.  We aim to use the spectra to confirm the low surface gravities of the objects in the near-infrared and to test the ability of the BT-Settl models to correctly handle the formation and gravitational settling of dust under reduced surface gravity conditions.  These spectra enrich the scarce sample of empirical near-infrared medium-resolution spectra of young late-type objects  beyond the M-L transition, especially for spectral type L3.  We describe in Section \\ref{observations} our observations and the associated data reduction. We present in Section \\ref{empirical_analysis} an empirical analysis of the spectral features in order to derive  near-infrared spectral types and confirm the young age of our targets. In Section \\ref{Spectralsynthesis} we describe the comparison of the atmospheric models to  the observed spectra. We discuss these comparisons  and derive updated target properties in Section \\ref{discussion}. ",
        "conclusions": "We obtained and analyzed seven VLT/ISAAC medium-resolution (R$\\sim$1500-1700) spectra of M9.5-L3 dwarfs classified at optical wavelengths and showing indications of low surface gravity. We built an age-sequence of M9.5 objects that allow us to pinpoint age-sensitive and gravity-sensitive features at medium-resolving powers. The comparison of our spectra to those of young reference brown dwarfs and companions, and of mature field dwarfs confirm that our objects have peculiar features in the near-infrared indicative  of low surface gravities and young ages. We also confirm the youth of our objects by calculating the equivalent widths of their KI lines and comparing these values per spectral type with the values obtained for young reference brown dwarfs and companions and  mature field dwarfs. We derived near-infrared spectral types based on dedicated spectral indices. These spectral types are in agreement with the optical classification, and confirm the coherence of the classification method.  The analysis revealed that the L2${\\gamma}$ object \\object{2MASS~J2322} provides a good match to the spectrum of the young planetary mass companion \\object{1RXS J160929.1-210524b}.    The spectra and SEDs of the objects can be reproduced by  the 2010 and 2013 BT-Settl atmospheric models.  The 2013 release of the models fits simultaneously the spectra and the SED for the same temperatures at all wavelengths. L0-L3$\\gamma$ dwarfs have nearly equal  temperatures around 1800 K. Nevertheless, we  identify that: \\begin{itemize} \\item the 2010 models do not  reproduce the 1.1-2.5 $\\mu$m spectral slope of some L2-L3 objects. \\item the H~band shape is not well reproduced by the BT-Settl 2013 models at solar metallicity. The problem disappears when new, but not as well-tested, models at super-solar metallicity are used, but these models remain mostly untested. \\end{itemize} Currently, all these discrepancies point out a lack of dust in the cloud models. The next version of the BT-Settl models will modify the treatment of the vertical mixing and of grain growth processes. These new models are expected to produce thicker clouds, and may solve the issues revealed by the ISAAC spectra.    The spectra of the objects will help to  confirm the membership of photometrically-selected candidates  in star-forming regions. Within the next few years, surveys on the next generation of planet imaging instruments such as SPHERE (Spectro-Polarimetric High-contrast Exoplanet REsearch) at VLT, GPI (Gemini Planet Imager) at Gemini South, ScEXAO (Subaru Coronagraphic Extreme AO Project) at Subaru, and LMIRCam (Large Binocular Telescope mid-infrared camera) at LBT should  provide a sample of a few dozen  young companions. Several planets similar to $\\beta$ Pictoris b should be unearthed and fall in the same temperature range as our objects. Therefore, our spectra will serve as precious benchmarks for the characterization of the physical and atmospheric properties of these companions."
    },
    "1402/1402.6192_arXiv.txt": {
        "abstract": "The gamma-ray-detected blazar 3C~454.3 exhibits dramatic flux and polarization variations in the optical and near-infrared bands. In December 2010, the object emitted a very bright outburst. We monitored it for approximately four years (including the 2010 outburst) by optical and near-infrared photopolarimetry. During the 2010 outburst, the object emitted two rapid, redder brightenings, at which the polarization degrees (PDs) in both bands increased significantly and the bands exhibited a frequency-dependent polarization. The observed frequency-dependent polarization leads us to propose that the polarization vector is composed of two vectors. Therefore, we separate the observed polarization vectors into short and long-term components that we attribute to the emissions of the rapid brightenings and the outburst that varied the timescale of days and months, respectively. The estimated PD of the short-term component is greater than the maximum observed PD and is close to the theoretical maximum PD. We constrain the bulk Lorentz factors and inclination angles between the jet axis and the line of sight from the estimated PDs. In this case, the inclination angle of the emitting region of short-term component from the first rapid brightening should be equal to 90$^{\\circ}$, because the estimated PD of the short-term component was approximately equal to the theoretical maximum PD. Thus, the Doppler factor at the emitting region of the first rapid brightening should be equal to the bulk Lorentz factor. ",
        "introduction": "Blazars are a type of active galactic nuclei with relativistic jets that are widely believed to be viewed at small angles to the line of sight. Blazars frequently show violent variations in flux and polarization, which vary on timescales ranging from minutes \\citep{Aharonian07,Sasada08} to years \\citep{Sillanpaa96}. On the timescale of months, blazars may turn out to be over 10 times brighter than in their quiescent state. These brightening phenomena are often called ``outburst''.  The blazar 3C~454.3 with a redshift $z=0.859$ \\citep{Jackson91} is one of the most famous blazar because it has emitted several large-amplitude outbursts. In 2005, it emitted a dramatic outburst that covered the range from radio to hard X-ray bands \\citep{Fuhrmann06,Pian06,Giommi06,Villata07}. The maximum brightness in the optical band reached $R=12.0$. Many authors reported that prominent outbursts occurred from 3C~454.3 in 2007, 2008 and 2009 and covered the range from radio to gamma-ray bands. Gamma-ray emission from the 2007 outburst was detected by the AGILE satellite \\citep{Vercellone08,Donnarumma09,Vercellone10}, and brightenings in other wavelengths were also detected \\citep{Villata08,Raiteri08a,Raiteri08b,Villata09a,Sasada10}. \\citet{Ghisellini07} and \\citet{Sikora08} proposed two possible explanations for the origin of seed photons for inverse Compton scattering emission in the GeV gamma-ray band. A large-amplitude gamma-ray brightening in the 2008 outburst was detected by the {\\it Fermi Gamma-ray Space Telescope} \\citep{Bonning09a,Abdo09}. Based on intensive monitoring, it was determined that the flux variations in the optical bands lagged no more than one day behind those of the GeV gamma-ray band \\citep{Vercellone09a,Donnarumma09,Bonning09a,Vercellone10}. In 2009, an outburst from 3C~454.3 occurred over all wavelengths: in the gamma-ray band it was detected by {\\it Fermi} and AGILE \\citep{Striani09a,Striani09b,Escande09,Striani10,Pacciani10,Ackermann10}, in the X-ray band it was detected by {\\it INTEGRAL} \\citep{Vercellone09b}, {\\it Swift}/XRT \\citep{Sakamoto09}, and {\\it Swift}/BAT \\citep{Krimm09}, in the optical bands it was detected by many groups \\citep{Villata09b,Bonning09b,Sasada09,Raiteri11,Sasada12}, and in the radio bands it was detected by \\citet{Raiteri11} and \\citet{Jorstad13}.  In September 2010, a prominent outburst from 3C~454.3 was detected in the GeV gamma-ray band by {\\it Fermi} and AGILE \\citep{Abdo11,Vercellone11}. The peak flux reached 85$\\times$10$^{-6}$ photons~cm$^{-2}$~s$^{-1}$ in December 2010. The optical continuum and emission-line fluxes increased simultaneously in this outburst \\citep{Vercellone11,Leon-Tavares13}, and the X-ray and radio fluxes also increased \\citep{Wehrle12,Jorstad13}.  In the present study, we report the results of using the Kanata telescope to monitor emission from 3C~454.3 using optical and near-infrared (NIR) photopolarimetry from July 2007 to January 2011: a period that includes the 2010 outburst. In December 2010, the polarization degree (PD) increased in the rapid brightenings of the 2010 outburst and the polarization difference between the optical and NIR bands became frequency dependent (this phenomenon is called ``frequency-dependent polarization'' or FDP). The temporal variations of polarization angle (PA) indicate that no rotation event occurred during this outburst. The paper is organized as follows: In section~2, we present the methods used for observation and analysis. In section~3, we report the results of the photometric and polarimetry observation. In section 4, we discuss the origin of the observed FDP and emission regions of the observed rapid brightenings and outburst. The conclusion is given in section 5. ",
        "conclusions": "We monitored blazar 3C~454.3 in the optical and near-infrared bands for approximately four years starting from July 2007, and we detected two types of variations in 2010: a large-amplitude outburst occurring on a timescale of $\\sim$months and extraordinary rapid brightenings occurring on a timescale of $\\sim$days. These brightenings had three features: (1) the $V-J$ bands reddened, (2) the degree of polarization (PD) increased in the $V$ and $J$ bands, and (3) the polarization was frequency dependent (FDP).  Based on these results, we suggest that the observed polarization vectors can be decomposed into two components, namely long- and short-term components. We separate the short-term component, estimating the long-term component for the rapid brightenings from the pre- and post-brightening data. The estimated PD of the short-term component for the rapid brightenings was higher than the observed maximum PD and was close to the theoretical maximum PD. This result indicates that the short-term emitting regions responsible for the rapid brightenings are different from the long-term emitting region that radiated the outburst. The blazar FDP gives us important constraints on $\\Gamma$, $\\Phi$, and $\\delta$. \\\\ \\\\ This work was supported by a Grant-in-Aid for JSPS Fellows."
    },
    "1402/1402.1856_arXiv.txt": {
        "abstract": "We show that the local type non-Gaussianity in a class of curvaton models is suppressed, i.e. the non-linearity parameters $f_{\\rm NL}$ and those related with higher order statistics can be at most $O(1)$, even if the curvaton energy density is subdominant at the decay. This situation is naturally realized in a very simple curvaton potential with quadratic term plus quartic term. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3581_arXiv.txt": {
        "abstract": "Mergers of compact stellar remnants are prime targets for the LIGO/Virgo gravitational wave detectors. The gravitational wave signals from these merger events can be used to study the mass and spin distribution of stellar remnants, and provide information about black hole horizons and the material properties of neutron stars. However, it has been suggested that degeneracies in the way that the star's mass and spin are imprinted in the waveforms may make it impossible to distinguish between black holes and neutron stars. Here we show that the precession of the orbital plane due to spin-orbit coupling breaks the mass-spin degeneracy, and allows us to distinguish between standard neutron stars and alternative possibilities, such as black holes or exotic neutron stars with large masses and spins. ",
        "introduction": "Compact stellar remnant mergers are the main targets of gravitational wave (GW) detectors such as advanced LIGO (aLIGO)~\\citep{Harry:2010zz} and advanced Virgo (Adv)~\\citep{Acernese:2007zze}, with predicted rates between a few and a few hundred per year at full design sensitivity~\\citep{Abadie:2010cf}. These systems take tens of minutes to sweep through the sensitive band of the detectors, entering the band at $\\sim10$Hz, and terminating in the kHz range with a violent merger lasting just a few milliseconds.  The final stages of the inspiral and merger proceed differently for black holes (BHs) and neutron stars (NSs), and in principle, this should allow us to identify the make-up of the system from the GW signal alone. However, the number of GW cycles in the signal and the aLIGO/AdV sensitivity fall off rapidly with increasing frequency, meaning that there is very little information past $\\sim 500$ Hz (less than 2\\% of the SNR).  Probes of BH physics and the equation of state of NSs will likely require multiple detections~\\citep{DelPozzo:2011pg, DelPozzo:2013ala}. An electromagnetic counterpart to the GW signal, such as a short-hard gamma-ray burst or an associated kilonova/macronova emission~\\citep{Metzger:2011bv}, would indicate that at least one of the bodies was a NS, but beaming effects or the luminosity of the signal may make detecting a counterpart difficult for the majority of mergers~\\citep{Abadie:2010cf,Aasi:2013wya}. Absent a counterpart, we must rely on the early inspiral to extract information about the make-up of the binary, which poses a challenge since finite size effects are completely negligible during inspiral~\\citep{Read:2009yp}. All we have to go on to decide the composition of the binary are the values of the masses and spins inferred from the inspiral signal.  General arguments based on stability and causality limit the mass and spin of NSs to the range $M\\in[0.1, 3.2 ]M_\\odot$ for the mass and $\\chi \\in[0,0.7]$ for the dimensionless spin magnitude,~$\\chi \\equiv |\\vec{S}|/M^{2}$, where $\\vec{S}$ is the spin angular momentum~\\citep{Rhoades:1974fn,Lattimer:2006xb,Yagi:2014bxa}. Realistic equations of state yield a tighter mass range $M\\in[1.0, 2.5] M_\\odot$. The observed range of masses and spins is somewhat tighter~\\citep{Lattimer:2006xb,Ozel:2012ax}: $M\\in[1.0, 2.0]M_\\odot$, $\\chi \\in[0,0.3]$. The old NSs that merge are expected to have spun down by magnetic breaking to the point where the maximum spin is much lower, $\\chi \\lesssim 0.05$, than in the general NS population~\\citep{Mandel:2009nx}. Furthermore, the standard isolated NSNS binary formation scenario ensures that after every common envelope phase (that tends to align the spins) follows a supernovae kick that misaligns the spins (unless the kick is in the orbital plane, though there is evidence that this is not the case~\\citep{Kaplan:2008qm}). Thus, we adopt the definition that {\\em normal} NSs seen by aLIGO/AdV have $M\\in[1, 2.5]M_\\odot$ and $\\chi \\leq  0.05$, and term NSs with larger masses or spins {\\em exotic}. Einstein's theory of gravity allows BHs to have spin in the range $\\chi \\in[0,1]$ with any mass. X-ray observations have identified stellar remnant BHs with $M\\in [3.6, 36]M_\\odot$ and $\\chi \\in[0,1]$. There is currently some debate as to the existence of a mass gap between NSs and BHs~\\citep{Ozel:2010su,Farr:2010tu,Belczynski:2011bn}, but for the purpose of determining whether a normal NS could be misidentified as a BH or an exotic NS, the existence of a gap is moot.  The early inspiral phase of a compact binary merger can be modeled analytically by expanding Einstein's equations in powers of the ratio of the orbital velocity to the speed of light, the so-called \\emph{post-Newtonian} (PN) approximation~\\citep{Blanchet:2013haa}. This ratio is small during the inspiral, with $v/c$ of $7\\%$ when the system enters the detector sensitivity band, reaching roughly $40\\%-60\\%$ by contact~\\citep{Bernuzzi:2014kca}. The PN approximation becomes less accurate as the system evolves through the band, eventually breaking down at the end of the inspiral phase. As all forms of energy couple to gravity, both the masses and spins leave an imprint on the binary orbit and the GWs emitted. The coupling between spin and orbital angular momentum can strongly affect the orbital trajectory and the GWs emitted in the inspiral phase.  The PN approximation can be used to construct a model of the GWs emitted during inspiral. The combination of such a GW model with a model for the instrument response yields templates for the signals as seen by the detector. Subtracting the model from the data produces a residual, and demanding that the residual is consistent with a model for the instrument noise defines a likelihood function. From this function and our prior knowledge we can derive a posterior distribution for the model parameters that are consistent with the observed data. It often happens that there are strong correlations between these parameters, limiting our ability to measure each parameter individually.  Recent work~\\citep{Hannam:2013uu} has suggested that the correlation between mass and spin~\\citep{Cutler:1992tc,Cutler:1994ys} may make it impossible to distinguish between a NSNS binary and a NSBH or a BHBH binary. This result hinges on a simplified waveform model that assumes that the spin and orbital angular momenta are perfectly aligned, and thus, spin-orbit induced precession~\\citep{springerlink:10.1007/BF00756587,Bohe:2012mr} is absent. However, we have no reason to expect the spin and orbital angular momenta to be aligned in stellar remnant binaries. Indeed, the NS binaries observed at much longer orbital periods are far from aligned and are precessing~\\citep{2002ApJ...576..942W,2005ApJ...624..906H,2008Sci...321..104B}. It has been hypothesized~\\citep{Hannam:2013uu,Baird:2012cu} that spin precession would not significantly alter the conclusions drawn using spin-aligned waveforms. We have tested this hypothesis and found, as first suggested by~\\citet{Cutler:1992tc}, that spin precession adds additional richness to the signals that almost completely breaks the mass-spin degeneracy,  producing an order-of-magnitude improvement in the extraction of the individual masses and spins, which allows us to distinguish between NSs and BHs. We show that normal NS binaries will not be mistaken for BHs or exotic NSs, but we cannot rule out the possibility that some exotic NSs or low mass/low spin BHs may be misidentified as normal NSs. ",
        "conclusions": "We showed that the inclusion of spin-precession in waveform templates breaks the degeneracy between the system's individual masses and spins, and allows us to distinguish between NSNS binaries and low-mass, small-spin NSBH or BHBH binaries. Moreover, even for signals with modest SNR, we can distinguish between ``normal'' and ``exotic'' NSs. These results open the door to population studies with the first GW detections, as well as coincident studies between the electromagnetic detection of short gamma-ray bursts and GWs. Indeed, if such a coincident observation is made, being able to identify the source from purely GW observations as a NS binary, a mixed binary or a BH binary would prove invaluable.  The results presented here are subject to several assumptions. First, the noise is assumed to be stationary and Gaussian, while in reality this may not be the case. Proper noise modeling along the lines described in~\\citet{Littenberg:2014oda,Cornish:2014kda} will help to restore performance to levels close to the ideal case. Second, calibration errors and non-stationary drifts in the noise spectrum should be marginalized over in a full analysis, but these mostly impact the amplitude parameters, and only have a small impact on the spin measurement. Third, the waveform model inaccuracies do not affect our estimates of the statistical errors at leading order~\\citep{Chatziioannou:2014bma}, so our conclusions will apply to more accurate waveform models.   \\newpage  \\emph"
    },
    "1402/1402.3062_arXiv.txt": {
        "abstract": "{The commonly used extinction laws of Cardelli et al. (1989) have limitations that, among other issues, hamper the determination of the effective temperatures of O and early B stars from optical and NIR photometry.} {We aim to develop a new family of extinction laws for 30 Doradus, check their general applicability within that region and elsewhere, and apply them to test the feasibility of using optical and NIR photometry to determine the effective temperature of OB stars.} {We use spectroscopy and NIR photometry from the VLT-FLAMES Tarantula Survey and optical photometry from HST/WFC3 of 30 Doradus and we analyze them with the software code CHORIZOS using different assumptions, such as the family of extinction laws.} {We derive a new family of optical and NIR extinction laws for 30 Doradus and confirm its applicability to extinguished Galactic O-type systems. We conclude that by using the new extinction laws it is possible to measure the effective temperatures of OB stars with moderate uncertainties and only a small bias, at least up to $\\ebv \\sim 1.5$ mag.} {} ",
        "introduction": "Astronomy is entering a time when massive photometric surveys allow us to obtain information about a very large number of objects. Projects such as Gaia and the Large Synoptic Survey Telescope (LSST) will reinforce this trend in the next decade. The main goal of these surveys is to measure the intrinsic properties of these objects, such as the effective temperature, luminosity, and metallicity of stars; the mass, age, and metallicity of stellar clusters; or the redshift and type of galaxies. These surveys include not only large numbers of targets, but also detailed calibration mechanisms that lead to (internal) precisions and (external) accuracies at the level of one hundredth of a magnitude. In other words, we have not only data in large quantities but also with high quality in the form of random and systematic errors that are significantly lower than what was typical twenty years ago. This high photometric quality is also extended to space missions such as the Hubble Space Telescope (HST) and is due to the stability of the space environment and the resources devoted to ensure the uniformity of the data.  Despite such high quality, there is (and always will be) one obstacle for the derivation of the intrinsic properties of astronomical objects: extinction. Every observation has to be corrected for the presence of dust between the target and the observer and that can be (and in many cases is) the main limitation. In the 1980s the great success of the International Ultraviolet Explorer (IUE) satellite prompted a revived interest in the subject of extinction that culminated with the groundbreaking work of \\citet[hereafter CCM]{Cardetal89}. that paper provided for the first time a family of extinction laws that extended from the IR to the UV while simultaneously characterizing the type of extinction with a single parameter, \\rv\\ (see \\citealt{Maiz13b} for a discussion on the name and the precise nature of the parameter). These two characteristics made the CCM laws a resounding success and the paper one of the most cited in astronomy in the last quarter of a century.  Despite their unquestioned relevance, different studies in the last two decades have revealed several issues with some aspects of the CCM laws:  \\begin{itemize} \\item the use of band-integrated [$E(B-V)$ and $R_V$] quantities to define the amount and type of extinction instead of their monochromatic equivalents (\\ebv\\, and \\rv, respectively)\\footnote{It cannot be emphasized enough that using $R_V$ to parameterize an extinction law is a serious mistake. $R_V\\equiv A_V/E(B-V)$ depends not only on the extinction law but also on the amount of extinction and the input SED. The reader is referred to Figure 3 of \\citet{Maiz13b} to quantify the effect. The parameter called $R_V$ in CCM is not really that (in the sense that an extinction law with a given value of that parameter does not yield that value of $A_V/E(B-V)$ for an arbitrary amount of extinction and an arbitrary SED), but a monochromatic value. In addition, this type of effect in broad-band photometry has been known at least since \\citet{Blan57} but appears to be overlooked by a significant fraction of the astronomical community.}; \\item the validity of a fixed extinction law in the NIR; \\item the functional form used in the optical; \\item the reality of the correlation between \\rv\\ and UV extinction; \\item the applicability of the laws beyond the \\ebv\\ and \\rv\\ values of the sample used to derive them; \\item the photometric calibration of the filters. \\end{itemize}  These issues are discussed in \\citet{Maiz13b}, where the reader is referred for details, and they are the reasons that prompted us to attempt an improvement of the CCM laws, concentrating on the correction for extinction for photometric data, their most commonly used application.  This paper is part of a series on the VLT-FLAMES Tarantula Survey. The reader is referred to the first paper, \\citet{Evanetal11a}, for details on the project. Within the series, this paper on the optical and NIR extinction law in 30 Doradus and its application to the determination of effective temperatures (\\teff) is part of a subseries on extinction and the ISM. The subseries started with the work of \\citet{vanLetal13} on diffuse interstellar bands and neutral sodium and will continue with another paper on the spatial distribution of extinction in 30 Doradus (Ma{\\'\\i}z Apell\\'aniz et al. in preparation).  We start by describing the spectroscopic and photometric data in this paper. We then perform different experiments with the data by processing them with CHORIZOS \\citep{Maiz04c}. The results are discussed and possible future work is described. The paper ends with three appendixes on (a) the detailed changes introduced by the new laws, (b) CHORIZOS and the spectral energy distributions (SEDs) used for this paper, and (c) the extinction along a sightline with more than one type of dust. ",
        "conclusions": ""
    },
    "1402/1402.1251_arXiv.txt": {
        "abstract": "{In this paper I provide a general framework based on $\\delta N$ formalism to study the features of unavoidable higher dimensional non-renormalizable K\\\"ahler operators for ${\\cal N}=1$ supergravity (SUGRA) during primordial inflation from the combined constraint on non-Gaussianity, sound speed and CMB dipolar asymmetry as obtained from the recent Planck data. In particular I study the nonlinear evolution of cosmological perturbations on large scales which enables us to compute the curvature perturbation, $\\zeta$, without solving the exact perturbed field equations. Further I compute the non-Gaussian parameters $f_{NL}$ , $\\tau_{NL}$ and $g_{NL}$ for local type of non-Gaussianities and CMB dipolar asymmetry parameter, $ A_{CMB}$, using the $\\delta N$ formalism for a generic class of sub-Planckian models induced by the Hubble-induced corrections for a minimal supersymmetric D-flat direction where inflation occurs at the point of inflection within the visible sector. Hence by using multi parameter scan I constrain the non-minimal couplings appearing in non-renormalizable K\\\"ahler operators within, ${\\cal O}(1)$, for the speed of sound, $0.02\\leq c_s\\leq 1$, and tensor to scalar, $10^{-22} \\leq r_{\\star} \\leq 0.12$. Finally applying all of these constraints I will fix the lower as well as the upper bound of the non-Gaussian parameters within, ${\\cal O}(1-5)\\leq f_{NL}\\leq 8.5$, ${\\cal O}(75-150)\\leq\\tau_{NL}\\leq 2800$ and ${\\cal O}(17.4-34.7)\\leq g_{NL}\\leq 648.2$, and CMB dipolar asymmetry parameter within the range, $0.05\\leq A_{CMB}\\leq 0.09$.}    \\begin{document} ",
        "introduction": "The primordial inflationary paradigm is a very rich idea to explain various aspects of the early universe, which creates the perturbations and the matter. For recent developments see Refs.~\\cite{Mazumdar:2010sa,Mazumdar:2011zd}. Usually inflation prefers slow-rolling of a single scalar field on a flat potential, which has unique predictions for the Cosmic Microwave Background (CMB) observables. The induced cosmological perturbations are generically random Gaussian in nature with a small tilt and running in the primordial spectrum which indicates that inflation must come to an end in our patch of the universe. But a big issue may crop up in model discrimination and also in the removal of the degeneracy of cosmological parameters obtained from CMB observations \\cite{Spergel:2006hy,Ade:2013uln,Ade:2013zuv,Ade:2013ydc}. Non-Gaussianity has emerged as a powerful observational tool to discriminate between different classes of inflationary models \\cite{Maldacena:2002vr,Bartolo:2004if,Komatsu:2009kd,Chen:2010xka,Komatsu:2010hc}. The Planck data show no significant evidence in favour of primordial non-Gaussianity, the current limits \\cite{Ade:2013ydc} are yet to achieve the high statistical accuracy expected from the single-field inflationary models and for this opportunities are galore for the detection of large non-Gaussianity from various types of inflationary models. To achieve this goal, apart from the huge success of cosmological linear perturbation theory, the general focus of the theoretical physicists has now shifted towards the study of nonlinear evolution of cosmological perturbations. Typically any types of nonlinearities are expected to be small; but, that can be estimated via non-Gaussian n-point correlations of cosmological perturbations. The so-called ``$\\delta N$ formalism'' (where $N$ being the number of e-foldings) \\cite{Starobinsky:1982,Salopek:1990,Sasaki:1995aw,Wands:2000dp,Lyth:2004gb,Lyth:2005fi,Mazumdar:2012jj,Sugiyama:2012tj} is a well accepted tool for computing non-linear evolution of cosmological perturbations on large scales ($k<<aH$), which is derived using the ``separate universe'' approach \\cite{Sasaki:1995aw,Wands:2000dp,Sasaki:1998ug,Sasaki:2007ay}. Particularly, it provides a fruitful technique to compute the expression for the curvature perturbation $\\zeta$ without explicitly solving the perturbed field equations from which the various local non-Gaussian parameters, $f_{NL}^{\\mathrm{local}},\\tau_{NL}^{\\mathrm{local}},g_{NL}^{\\mathrm{local}}$ and CMB dipolar asymmetry parameter \\cite{Erickcek:2008sm,Lyth:2013vha,Wang:2013lda,Dai:2013kfa,Mazumdar:2013yta}, $A_{CMB}$ are easily computable~\\footnote{One can also compute all the non-gaussian parameters, $f_{NL}^{\\mathrm{local}},\\tau_{NL}^{\\mathrm{local}},g_{NL}^{\\mathrm{local}}$ and CMB dipolar asymmetry parameter, $A_{CMB}$ using {\\it In-In formalism} in the quntum regime. But the inflationary dynamics responsible for the interactions between the modes occurs at the super-horizon scales within the effective theory setup proposed in this paper. Here I use $\\delta N$ formalism as-(1) it perfectly holds good at the super-horizon scales and (2) are also independent of any kind of intrinsic non-Gaussianities generated at the scale of horizon crossing.}.    We will be using the following constraints on the amplitude of the power spectrum, $P_s$, spectral tilt, $n_s$, tensor-to-scalar ratio, $r$, sound speed, $c_{s}$, local type of non-Gaussianity, $f_{NL}^{\\mathrm{local}}$ and $\\tau_{NL}^{\\mathrm{local}}$ and CMB dipolar asymmetry from Planck data throughout the paper \\cite{Ade:2013uln,Ade:2013zuv,Ade:2013ydc,Ade:2013nlj}:  \\begin{eqnarray} \\label{pow} \\displaystyle \\ln(10^{10}P_{s})=3.089^{+0.024}_{-0.027}~~(~at~2\\sigma~~ CL)\\,, \\\\ \\label{ns} \\displaystyle  n_{s}=0.9603\\pm 0.0073~~(~at~2\\sigma~~ CL)\\,, \\\\ \\label{rten} \\displaystyle r\\leq 0.12~~(~at~2\\sigma~~ CL)\\,, \\\\ \\label{f-nl} \\displaystyle 0.02 \\leq c_s \\leq 1~~(~at~2\\sigma~~ CL)\\,, \\\\ \\label{f-nl1}\\displaystyle f_{NL}^{\\mathrm{local}}=2.7\\pm 5.8~~(~at~1\\sigma~~ CL)\\,, \\\\ \\label{f-nl2}\\displaystyle \\tau_{NL}^{\\mathrm{local}}\\leq2800~~(~at~2\\sigma~~ CL)\\,, \\\\ \\label{f-nl3}\\displaystyle A_{CMB}=0.07\\pm 0.02~~(~at~2\\sigma~~ CL). \\end{eqnarray}   In this paper I will concentrate our study for Hubble induced inflection point MSSM inflation derived from various higher dimensional Planck scale suppressed non-minimal K\\\"ahler operators in ${\\cal N}=1$ supergravity (SUGRA) which satisfies the observable universe, and it is well motivated for providing an example of visible sector inflation.   In section \\ref{setup}, I will briefly review the setup with one heavy and one light superfield which are coupled via non-minimal interactions through K\\\"ahler potential. In section \\ref{N-M} I discuss very briefly the role of various types of Planck suppressed non-minimal K\\\"ahler corrections to model a Hubble induced MSSM inflation for any D-flat directions. Hence in section \\ref{BNV} I present a quantitaive analysis to compute the expression for the local types of non-Gaussianin parameters and CMB dipolar asymmetry parameter which characterize the bispectrum and trispectrum using the $\\delta N$ formalism. For the numerical estimations I analyze the results in the context of two D-flat direction, $\\widetilde L\\widetilde L\\widetilde e$ and $\\widetilde u\\widetilde d\\widetilde d$ within the framework of MSSM inflation \\cite{Allahverdi:2006iq,Allahverdi:2006cx}. ",
        "conclusions": "In this paper, I have shown that in any general class of ${\\cal N}=1$ SUGRA inflationary framework, the behaviour of K\\\"ahler potential in presence of non-minimal K\\\"ahler corrections in effective theory setup are constrained via the non-minimal couplings of the non-renormalizable gauge invariant K\\\"ahler higher dimensional Planck scale suppressed operators from the observational constraint on non-Gaussianity, sound speed and CMB dipolar asymmetry as obtained from the Planck data. In the present setup the hidden sector based heavy field is settled down in its potential via its Hubble induced vacuum energy density. In particular, for the numerical estimations in this paper I have used a very particular kind of (inflection point) inflationary model, which is fully embedded within MSSM, where the inflaton is made up of $\\widetilde L\\widetilde L\\widetilde e$ and $\\widetilde u\\widetilde d\\widetilde d$ gauge invariant D-flat directions. However the prescribed methodology holds good for other kinds of inflationary models too.  Further I have scanned the multiparameter region characterized by the Hubble induced mass parameter, $c_{H}$, A-term, $a_{H}$ and the scale of the heavy field $M_{s}$, where I have satisfied the current Planck observational constraints on the, inflationary parameters: $P_S,~n_S,~c_s,r_*$ (within $2\\sigma$~CL), non-Gaussian parameters: $f_{NL}^{local},\\tau_{NL}^{local}$ (within $1\\sigma-1.5\\sigma$~CL) and CMB dipolar asymmetry parameter $A_{CMB}$ (within $2\\sigma$~CL). Consequently the non-minimal K\\\"ahler couplings, $a,b,c,d$ are fixed within $\\sim {\\cal O}(1)$ in the present effective theory setup. Finally, using this methodology, I have obtained the theoretical upper and lower bound on the non-Gaussian parameters within the range, ${\\cal O}(1-5)\\leq f_{NL}\\leq8.5$, ${\\cal O}(75-150)\\leq\\tau_{NL}<2800$ and ${\\cal O}(17.4-34.7)\\leq g_{NL}\\leq 648.2$, and the CMB dipolar asymmetry parameter within, $0.05\\leq A_{CMB}\\leq0.09$, which satisfy the observational constraints stated in Eq~(\\ref{pow}-\\ref{f-nl3}), as obtained from Planck data.  There is also a future prospect of upgrading the present methodology proposed in this paper by studying the further stringent phenomenological constraints on the non-minimal couplings $a,b,c,d$, as appearing in the context of higher dimensional Planck scale suppressed K\\\"ahler operators within ${\\cal N}=1$ SUGRA by imposing the constraint on Higgs mass \\cite{Aad:2012tfa,Chatrchyan:2013lba} and the dark matter abundance \\cite{Ade:2013zuv,Boehm:2012rh} via WIMPy baryogenesis scenario (see also Refs.~\\cite{Cui:2011ab,Bernal:2012gv,Bernal:2013bga} for the various theoretical issues). A detailed analysis on these aspects have been reported shortly as a separate paper \\cite{Sayan:2014}."
    },
    "1402/1402.5969_arXiv.txt": {
        "abstract": "{Several recent studies have found that planet migration in adiabatic disks differs significantly from migration in isothermal disks. Depending on the thermodynamic conditions, that is, the effectiveness of radiative cooling, and on the radial surface density profile, planets migrate inward or outward. Clearly, this will influence the semimajor axis-to-mass distribution of planets predicted by population-synthesis simulations.} {Our goal is to study the global effects of radiative cooling, viscous torque desaturation, gap opening and stellar irradiation on the tidal migration of a synthetic planet population.} {We combined results from several analytical studies and 3D hydrodynamic simulations in a new semi-analytical migration model for the application in our planet population synthesis calculations.} {We find a good agreement of our model with torques obtained in 3D radiative hydrodynamic simulations. A typical disk has three convergence zones to which migrating planets move from the in- and outside. This strongly affects the migration behavior of low-mass planets. Interestingly, this leads to a slow type II like migration behavior for low-mass planets captured in these zones even without an ad hoc migration rate reduction factor or a yet-to-be-defined halting mechanism. This means that the new prescription of migration that includes nonisothermal effects makes the previously widely used artificial migration rate reduction factor obsolete.} {Outward migration in parts of a disk helps some planets to survive long enough to become massive. The convergence zones lead to potentially observable accumulations of low-mass planets at certain semimajor axes. Our results indicate that more studies of the mass at which the corotation torque saturates are needed since its value has a main impact on the properties of planet populations. } ",
        "introduction": "\\label{sect:introduction} The huge diversity found in the properties of extrasolar planets is challenging to reproduce with global theoretical planet formation models. The goal of such a model is to explain all the different planet types, which range from  low-mass rocky planets such as Kepler-10 b \\citep{batalha2011} and multiplanet systems like our solar system to high-mass planets orbiting far from their star, such as NR 8977 \\citep{Maroisetal2008}.  The only way to study this problem is to use the results of global formation and evolution models and to compare them statistically with the steadily increasing number of known planets and their physical properties. This is done in planet populations synthesis calculations, in which the evolution of one or several planets and the harboring protoplanetary disk is calculated at the same time in Monte Carlo simulations to create whole populations of planets. Several groups presented studies based on this method, for example, \\citet{idalin2004},\\citet{idalin2010}, \\citet{Thommesetal2008},\\citet{Miguelbrunini2008},\\citet{HellaryNelson2012}, or by our group, \\citet{mordasinialibert2009a},\\citet{mordasinialibert2009b},\\citet{mordasini2012a},\\citet{mordasini2012b}, and \\citet{alibertmordasini2011}.  One general result found in all these models is that giant planets close to the star (``hot'' Jupiters) do not form \\textit{insitu}: the extrapolation of disk properties found at larger distances to small distances indicates that there is probably not enough solid material close-in to form a sufficiently large core that whould be able to accrete gas. The amount of material a core can accrete locally is given by the isolation mass $\\miso$. According to the empirical minimum-mass solar nebula model (MMSN), the isolation mass is only a fraction of the Earth mass ($\\mearth$) inside of 1 AU \\citep{idalin2004}. Therefore an increase of solid matter by roughly two orders of magnitude compared with the MMSN would be needed \\citep{idalin2004}. This means that to explain the close-in ``hot'' Jupiters, they would have to form initially at larger separations from the star and move inward by some mechanism (e.g., planet-planet scattering \\citep{RasioFord1996}, Kozai mechanism \\citep{NagasawaIda2008}, or tidal interactions with the gas disk \\citep{goldreichetremaine1980,tanakatakeuchi2002}). In the planet formation model used in this work, we consider only one core per disk, which means that only migration caused by tidal interactions can be studied. This migration is generally described by two different regimes that depend on the mass of the protoplanet. The first is type I migration for low-mass planets, which are too small to form a gap in the disk, and the second is type II migration for planets that open a gap \\citep{DangeloHenning2002}.  In our previous work \\citep[e.g. ,][]{mordasinialibert2009a,alibertmordasini2011}, we used the results obtained for isothermal disks reported in \\citet{tanakatakeuchi2002} to calculate type I migration rates. The migration rate in this model only depends on the disk surface density profile and not on the temperature profile of the disk because it assumes a globally isothermal disk. Migration rates obtained with this model always lead to rapid inward migration in disks with profiles similar to the MMSN. We showed \\citep{mordasinialibert2009a} that to obtain a synthetic population compatible with observations, one needs to artificially reduce the isothermal type I migration rate by a large factor. \\citet{idalin2004} used a similar type I migration prescription and found necessary reduction factors of $\\lesssim0.1$ for the migration rate.  Recent studies of type I migration in 2D or 3D hydrodynamical simulations also found outward migration for some masses or semimajor axes, depending on the disk temperature structure \\citep{massetdangelo2006,paardekoopermellema2008,kleybitsch2009}. More analytical work derived a formulation that could be used in planet population synthesis calculations \\citep{casolimasset2009,massetcasoli2009,massetcasoli2010}. Finally, \\citet{paardekooperbaruteau2010} derived a formalism for type I migration in adiabatic or locally isothermal disks and improved this even more in \\citet{paardekooper2011}.  In the present work we describe a semi-analytical type I migration model that can be applied to a wider range of planet and disk properties than that of \\citet{paardekooper2011}. For this we used the adiabatic and locally isothermal migration equations from \\citet{paardekooperbaruteau2010} and ratios of relevant time scales to determine the transition between different regimes. We then studied the global consequences of the physics included in the migration model in new sets of population synthesis calculations.  As an overview of this work, the new semi-analytical model we created is introduced in Section 2, where we also compare torques obtained with this model with torques obtained with the model of \\citet{paardekooper2011} and data obtained in 3D radiative hydrodynamic simulations of \\citet{kleybitsch2009} and \\citet{bitschkley2011}. We discuss in Section 3 a reference synthetic planet population calculated with the nominal model, and in Section 4 we study the effect of different migration models on planetary populations. In Section 5 we draw our conclusions and summarize our results. ",
        "conclusions": "We have compiled a prescription for type I migration based on the latest hydrodynamic simulations of planet-disk interactions. We tested the influence of the prescription on the outcome of the planetary population synthesis calculations. Our migration model is based on the combination of results from \\citet{paardekooperbaruteau2010} and different time scales to distinguish between different migration regimes. These time scales reflect the thermodynamical behavior of the interaction between the planet and the surrounding disk like the horseshoe drag.  We first compared the migration torques of this model with a model of \\cite{paardekooper2011} and with radiative-hydrodynamical simulations from \\citet{kleybitsch2009} and \\citet{bitschkley2011}. The comparison of the theoretical torque curves with data from  \\citet{kleybitsch2009} and \\citet{bitschkley2011} suggests that with an adjustment of the viscous time scale in the calculation of the saturation mass (the mass when the corotation torque starts to vanish) by a factor of $\\fvisc \\approx 0.55$ our model reproduces the torque better.  We also showed the global effects of different parameters of various migration models in a number of population synthesis calculations. Here, the comparison of our nominal BMF model ($\\fvisc \\approx 0.55$) and the migration model of \\citet{paardekooper2011} indicates similar results (Sect. \\ref{satmasssyn}) even with the difference shown in the torque curves.  Like \\citet{lyrapaardekooper2010},\\citet{KretkeLin2012}, and \\citet{HellaryNelson2012}, we also find with our prescription that nonisothermal type I migration leads to  convergence zones (CZ), that is, points in a disk to which planets migrate to from the inside and outside. As in previous studies \\citep{HellaryNelson2012}, we find that the convergence zones move inward as the disk evolves and take the captured planets with it. This occurs on the slower time scale of disk evolution and therefore the captured planets are trapped in it and also only migrate on this time scale. The planets leave this zone when their horseshoe drag saturates.  The difference between the migration rate of planets captured in a CZ and the migration rate in the saturated regime is significant. This means that an increase of the mass where saturation occurs by a modest factor of 2, for example, the time spent in rapid-saturated type I migration is significantly shortened and therefore also the extent of migration. The level of saturation at a given mass, and the mass at which saturation begins, are among the critical aspects for the evolution of a giant planet since small changes by a factor of 2 in $\\fvisc$ can change the final distribution in the mass semimajor axis diagram by a measurable degree. But a similar degree of change in the final semimajor axis-mass distribution was seen when we changed the description of the Lindblad torque to the formula found by  \\citet{massetcasoli2010}.  Finally, we determined the formation tracks of a planet, illustrating that under certain conditions, a planet can loose almost all of its gas mass during its evolution. This mass loss can lead to a stop of migration for a few $10^5$ years when the mass loss is strong enough to desaturate the horseshoe drag. The reason for the envelope mass loss is a jump in solid accretion rate, which is caused by the migration of the planet from a solid-depleted into a solid-rich region of the disk. This behavior strongly depends on the treatment of the planetesimal disk. Here we did not evolve the disk or changed the random velocities because of a giant planet. Both change the accretion rate and therefore whether or how this mass loss occurs.  Neither did we consider random variations in the torque due to turbulent density variations in the disk. Recent studies showed that in some cases random migration due to turbulence can dominate the migration behavior for low-mass planets \\citep{PapaloizouNelson2004,UribeKlahr2011}. Depending on the strength of the random torque, it could push planets from the inner into the outer convergence zone even for the stronger Lindblad torque of \\citet{paardekooper2011}.  We made no detailed comparison with observed extrasolar planet populations because the migration model is only a small part of the improvements to the overall model, and defer such comparisons to future work. But we quantified the impact of different model settings by studying the fractional yield of different planet types in a synthesis.  With the new nonisothermal migration model described here or that of \\citet{paardekooper2011} there are still about $50\\%$ of all low-mass planets lost in the innermost part of the disk. However, it is about a factor two better than the isothermal migration model without any artificial factors. One way to reduce the number of ``lost'' planets is an increase of the saturation mass. A lower critical mass for runaway gas accretion can also help due to the faster transition into type II migration \\citep{HoriIkoma2011}. Furthermore, there are hints that there are additional effects leading to outward migration of planets in 3D simulations of magneto-hydrodynamic disks \\citep{UribeKlahr2011}.  We also note that up to now there is a shortcoming of all analytic torque predictions: they all neglect the fact that the horseshoe region is over-wide compared with the prediction used here when $q/h^3 \\approx 1$ (\\citet{massetdangelo2006}). The fast growth of the width of the horseshoe region in that mass range (a few ten Earth masses to a hundred Earth masses) yields a boost of the corotation torque, which is a strong effect. In the same mass range gap opening and the transition into type II migration occur. And as shown in Section \\ref{compmodel}, saturation plays a main role in the change of the direction of the torque in that mass range as well. Moreover, our fit of the torques data from  \\citet{kleybitsch2009} and \\citet{bitschkley2011} depends on the progression of the torque curve in that mass range. This means that even when our fit produced the torque data well, a closer study of the torque in the mass range might uncover new effects, that may change the outcome of population synthesis calculations.  Nevertheless, the CZ or a similar effect might explain the concentration of planets in clusters one finds in distance-mass diagrams of observed extrasolar planets for high- and low-mass planets \\citep{Mayor2011}. In particular, the absence of close-in, low-mass planets in \\cite{mordasinialibert2009b} is not seen any longer with the updated migration model.  Additional comparisons of our results with new data from radial velocity measurements and Kepler data will be important, especially when we combine the new migration model with new improvements of our model, i.e., the long term evolution of giant planets \\citep{mordasini2012a} and the concurring evolution of multiple planets per disk \\citep{idalin2010,Alibert2013}. Multiple cores in one disk might be collected into one CZ and form one larger core. Therefore the convergence region may function as a large feeding zone of solid matter onto a core captured in a convergence zone and thus enhance the solid accretion rate \\citep{HellaryNelson2012, Sandor2011, Horn2012}. Moreover, in particular the low-mass planets of the inner cluster only migrated through the inner convergence region. The outer boundary of this region is due to the transition in the gas opacity at the iceline. This means that all these planets moved only through the ice-free part of the protoplanetary disk. We will also study the global effects of different viscous $\\alpha$ values and deadzones on the behavior of the convergences zones with population synthesis calculations \\citep[cf.][]{HasegawaPudritz2011}. This brings us closer to a description of orbital migration without ad hoc efficiency factors.   \\addcontentsline{toc}{chapter}{Bibliography}"
    },
    "1402/1402.0899_arXiv.txt": {
        "abstract": "We present the first systematic physical modelling of the time-lag spectra between the soft (0.3--1 keV) and the hard (1.5--4 keV) X-ray energy bands, as a function of Fourier frequency, in a sample of 12 active galactic nuclei which have been observed by \\textit{XMM}-\\textit{Newton}. We concentrate particularly on the negative X-ray time-lags (typically seen above $10^{-4}$ Hz) i.e.\\ soft band variations lag the hard band variations, and we assume that they are produced by reprocessing and reflection by the accretion disc within a lamp-post X-ray source geometry. We also assume that the response of the accretion disc, in the soft X-ray bands, is adequately described by the response in the neutral \\fa line at 6.4 keV for which we use fully general relativistic ray-tracing simulations to determine its time evolution. These response functions, and thus the corresponding time-lag spectra, yield much more realistic results than the commonly-used, but erroneous, top-hat models. Additionally we parametrize the positive part of the time-lag spectra (typically seen below $10^{-4}$ Hz) by a power-law. We find that the best-fitting BH masses, $M$, agree quite well with those derived by other methods, thus providing us with a new tool for BH mass determination. We find no evidence for any correlation between $M$ and the BH spin parameter, $\\alpha$, the viewing angle, $\\theta$, or the height of the X-ray source above the disc, $h$. Also on average, the X-ray source lies only around 3.7 gravitational radii above the accretion disc and $\\theta$ is distributed uniformly between 20 and 60\\degr. Finally, there is a tentative indication that the distribution of $\\alpha$ may be bimodal above and below 0.62. ",
        "introduction": "\\label{sect:intro} In the current paradigm, active galactic nuclei (AGN) contain a central black hole (BH) which is fed, in most cases, by an optically thick and geometrically thin accretion disc as a result of matter transportation inwards and angular momentum outwards. This disc radiates as a series of black-body components \\citep{shakura73} with peak emission at optical-ultraviolet wavelengths \\citep{malkan83}. Part of this radiation is assumed to be Compton up-scattered within a mildly relativistic hot electron cloud, often called the \\textit{X-ray corona}. This medium is often approximated as a point source (representing the centroid of the X-ray emitting source), lying above the central BH, on the axis of symmetry of the system (i.e.\\ BH spin axis). This arrangement is known as the `\\textit{lamp-post geometry}'.\\par The photons, from the X-ray source, form a power-law spectrum, and depending on the location of the X-ray source, a substantial part of them may illuminate the accretion disc. In this case, they are either Compton scatted by free or bound electrons, or photoelectrically absorbed followed by fluorescent line emission, or by Auger de-excitation. This yields the so-called `reflection spectrum' consisting of a number of emission and absorption lines (mainly below 1 keV), together with the 6.4 keV \\fa emission line from neutral material, which is the strongest X-ray spectral feature \\citep{george91}.\\par Depending on the proximity of the reflection process to the BH, the various lines can undergo relativistic broadening \\citep{fabian89,laor91}. These relativistically broadened lines may account for the observed \\textit{soft X-ray excess} \\citep{crummy06}, where the X-ray spectrum below 1 keV (soft band) lie above the power-law extrapolation of the continuum, usually measured in the 1.5--4 keV (hard band). However, alternative interpretations for the soft excess are also possible e.g.\\ Comptonisation of disc photons from the X-ray source \\citep{page04}, Comptonisation of disc photons within a hot layer on the accretion disc \\citep{wang03,done12}.\\par In addition to the primary and the reflected components, several AGN exhibit absorption features in their X-ray spectra as well. These feature can arise by: i) outflows which are either thermally or radiatively or magnetically driven winds \\citep{cappi06} and can reach mildly relativistic speeds \\citep{tombesi10}, ii) inflows \\citep{krug10}, and iii) X-ray absorption clouds \\citep{risaliti02}. The position of these structures varies from tens up to thousands of gravitational radii\\footnote{For a BH mass of $2\\times10^{6}$ \\ms\\ photons need $t_{\\rm g}=r_{\\rm g}/c=9.8$ s to travel a distance of 1 \\rg.}, $r_{\\rm g}=G M_{\\rm BH}/c^2$.\\par In this framework, and taking advantage of the highly variable behaviour of AGN, detection of time delays, as a function of the Fourier frequency, between the soft band and the hard band photons, can shed light on the X-ray emission mechanism and the geometry of these systems. Currently the observed negative X-ray time-lags (i.e.\\ soft-band variations lagging the hard-band variations; NXTL hereafter), has triggered a great deal of scientific interest on interpreting their nature. After their first tentative detection in the AGN Ark\\;564 \\citep{mchardy07}, where an origin in reflection from the accretion disc was first proposed, the first statistical significant detection came from \\citet{fabian09} for the AGN 1H\\;0707-495. Then, \\citet{emmanoulopoulos11b} found that this time delayed X-ray behaviour is much more common than was initially thought by analysing the data sets from two widely studied bright AGN, MCG--6-30-15 and Mrk\\;766. Subsequently, \\citet{demarco13} in a systematic analysis of 32 AGN, found a total of 15 AGN  exhibiting NXTLs with high statistical confidence.\\par The opposite time delayed behaviour i.e.\\ positive X-ray time delays (hard-band variations lag the soft-band variations), has been known for quite some time in both AGN \\citep[e.g.][]{papadakis01,mchardy04,arevalo06,sriram09} and X-ray binaries \\citep[XRBs; e.g.][]{miyamoto89,nowak96,nowak99}. Although positive time-lags are expected in the standard Comptonisation process within the X-ray source \\citep{nowak99}, they can also be produced by diffusive propagation of perturbations in the accretion flow \\citep{kotov01}.  \\par The origin of NXTL still remains unclear. Assuming a lamp-post geometry, in which the reflection of the hard X-rays occurs very close to the BH (few \\rg), the NXTL arise from the difference in the path-lengths between the soft (reflected photons) and the hard X-ray photons (coming directly from the X-ray source) to the observer \\citep{fabian09,zoghbi10,zoghbi11,cackett13,fabian13}. On the other hand  based on the large-scale distant scattering scenario \\citep{miller10_1h0707,legg12}, NXTLs might be caused by scattering of X-rays as they pass through, or are scattered from, a distant (tens up to hundred \\rg) absorbing clumpy medium (e.g.\\ wind, outflow, cloud), that partially covers the X-ray source, and whose opacity decreases with increasing energy.\\par Currently modelling of NXTLs is usually done in terms of simple top-hat impulse response functions \\citep[THIRF, hereafter;][]{miller10_1h0707,zoghbi11,emmanoulopoulos11b}. This approach is, however, just a parametrisation of the NXTL spectra and does not carry any physical information about the geometry and the physical properties of the BH (i.e.\\ mass, spin). In order to properly model the impulse response function, one has to take into account in detail the various general relativistic (GR) effects which affect the geometric path of both the hard and reflected X-rays. A first attempt at such a modelling of NXTLs was performed by \\citet{chainakun12} for the AGN 1H\\;0707-495 using a GR light-bending model and a moving X-ray source in order to take account of the observed X-ray variability. Their conclusion was that a more complex physical model is required in order to explain both the source's geometry and intrinsic variability. More recently, \\citet{wilkins13} using the full GR treatment with a variety of different geometries of corona, for the same AGN 1H\\;0707-495, found the the X-ray source  extends radially outwards to around 35 \\rg\\ and at a height of around 2 \\rg\\ above the plane of the accretion disc and that propagating fluctuations might account for the positive low frequency time delays.\\par In this paper we perform the first systematic analysis of the time-lag spectra of a sample of 12 AGN using fully general relativistic impulse response functions (GRIRFs, hereafter). These functions are generated using a lamp-post model with variable BH mass, BH spin parameter, viewing angle and height of the X-ray source above the disc. The paper is organised as follows: Initially, in Section~\\ref{sect:obs} we present the observations and data reduction procedures. Then, in Section~\\ref{sect:lampPost_irf} we describe the lamp-post model and the method that we use to derive the GRIRFs. In the next section, we outline the procedure for the estimation of the time-lag spectra models from the corresponding GRIRFs, and in Section~\\ref{sect:fits} we describe the fitting methodology. Section~\\ref{sect:res} contains the results of the best-fitting time-lag spectral models and finally in the last section we give a summary of our results and we discuss our conclusions. Throughout the paper the error estimates for the various physical parameters correspond to the 68.3 per cent confidence intervals unless otherwise stated. Similarly, the error bars of the plot points in all the figures indicate the 68.3 per cent confidence intervals. ",
        "conclusions": "\\label{sect:sum_disc} We have modelled the time-lag spectra of 12 AGN, using with the highest-quality \\textit{XMM}-\\textit{Newton} observations, in terms of signal-to-noise ratio and observational length. We have employed physically realistic GRIRFs for the case of the lamp-post geometrical model in which the X-ray source lies above the BH. This is the first time that such a fully GR approach is employed to model in a statistically robust way the time-lags spectra of a large number of local AGN. Our results imply that such a model can adequately describe the observed X-ray time delays and thus it is a realistic representation of the innermost regions of most AGN.\\par We find that the observed time-lag spectra of all objects are fully consistent with the hypothesis of X-ray reverberation from the innermost part of the accretion disc. The main results from our analysis are summarised below. \\begin{itemize} \\item The best-fitting BH-masses are in very good accordance with those derived from other, independent, methods, e.g.\\ optical reverberation mapping. \\item The best-fitting viewing angles imply an opening angle of around $80\\degr$ for the putative molecular torus, which is fully consistent with current estimates for the ratio of Type II over Type I Seyfert galaxies in the local Universe. \\item There is no correlation between the BH mass and any of the lamp-post model parameters, i.e.\\ spin, viewing angle and height (Fig.~\\ref{fig:lampPostCorrel}). \\item There is a tentative evidence for bimodality in the distribution of the best-fitting spin parameters above and below $\\alpha=0.62$: one group with a low BH-spin (with a mean of 0.35) and another group with a high BH spin (with a mean of 0.84), respectively. In any case, our results indicate, with a high significance, that the BH spin for member of the former group is not consistent with 1, while the BH spin for the objects in the latter group is not consistent with 0. But we do not find any correlation between BH mass and BH spin: our results suggest the existence of both Schwarzschild and Kerr BHs, at any given BH mass. \\item The average X-ray source height is 3.7 \\rg, with little dispersion. As these heights are an approximation to the X-ray source size, they imply a very small source size. \\end{itemize} Our modelling is limited by a few assumptions, which can potentially affect the validity of our results. To start with, as we have mentioned in Section~\\ref{ssect:relativ_respFunct}, strictly speaking, our GRIRFs do not correspond to the actual response of the soft band but rather to the response of the neutral fluorescent \\fa line, at 6.4 keV (in the rest frame of the accretion disc). However, if both the \\fa line and the soft band photons are generated in the same part of the accretion disc, there should not be dramatic differences between the two GRIRFs.\\par Perhaps, more important may be our assumption that the contribution of the reflection spectrum, in the soft X-ray band is equal to 30 per cent of the total observed flux for all objects (see also Appendix~\\ref{app:reflect_fracti}). If the contribution of the continuum emission in the soft band were insignificant (in which case $k$ in equation~\\ref{eq:soft_emis} is zero by definition), then the value of $f$ would not affect our results at all, since the time-lag spectrum, at a given frequency, would be determined only by the response function, $\\Psi(t)$ (equation~\\ref{eq:timeLagSpectra}). In the opposite case, when $k\\neq0$ and $f<<1$, the cross-correlation between the soft and hard band is basically equal to the autocorrelation of the continuum and thus the time-lags will be equal to those of the continuum (if any e.g.\\ positive time lags). Intermediate values of $f$ predict time-lag values between those of the continuum, and those due to the reprocessing component only. As we show in Appendix~\\ref{app: reflect_fracti} the resulting time-lag spectra depend quite a lot on the value of $f$. However, as we explained at the beginning of Section~\\ref{sect:fits}, we could not allow $f$ to be a variable parameter during the model fitting, due to the limited number of points in the observed time-lag spectra. However, the value of $f$ which we choose i.e.\\ 0.3, is the best current estimate based on detailed X-ray spectral analysis \\citep{crummy06}.\\par For similar reasons, to neglect the presence of a reflection component in the 2-\u20134 keV energy band, which we considered as a proxy for the direct continuum emission, will also affect our results. This issue has been discussed by \\citet{wilkins13}, who show that the presence of a reflection component in the continuum hard X-ray band does affect the strength of the reflection component in the resulting time-lag spectra. However, full investigation of this issue would have been computationally prohibitive in our case, where we try to model the time-lags spectra of a large number of AGN. Nevertheless, we believe that the good agreement between the BH mass estimates, from the present work, and the estimates resulting from other independent methods, as well as the reasonable estimates that we get for the viewing angles for all the objects in the sample, imply that our results should be a valid representation of the true values of the other model parameters i.e.\\ BH spin and X-ray source height and size as well. This could possibly also indicate that, based on the current data quality, we are not able to discern significant differences in the fitting parameters between the actual response of the soft band and that of the neutral \\fa line.\\par We note in passing that, given the good agreement between the BH estimates presented in this work, and others from the literature, the time-lag spectrum modelling with physically motivated models, is yet another way of measuring the BH mass in these systems. The uncertainty of our BH mass estimates is rather large, but this is not due to the method itself, but rather to the quality of the time-lags spectra, even though we have used the longest X-ray data sets currently available, obtained by \\textit{XMM}-\\textit{Newton}. In the case of objects with a BH mass larger than few times $10^7$ \\ms, the main limitation arises from the duration of the available X-ray light curves. For these objects, all the largest amplitudes time-lag spectral features (e.g.\\ the the most negative time-lag estimate) are shifted towards frequencies which are lower than the lowest accessible frequencies, defined by the length the light curves, which consist, in the best cases, of around 120 ks, or few times $10^{-5}$ Hz in the frequency domain. As a result, the only time-lag spectral features available to fit are of small amplitudes, resulting in larger BH mass fitting uncertainties (i.e.\\ the $\\chi^2$ regions around the best fitting-values are characterised by a smaller curvature and thus larger errors).\\par From the other hand, for objects with smaller BH mass, below $10^6$ \\ms, the largest time-lag amplitude variations are shifted towards the high Fourier frequencies above few times $10^{-3}$ Hz. In these frequencies, due to the Poisson noise, the very low coherence of our light curves prohibits any meaningful  physical modeling. Thus, in these cases we need to sample thoroughly the small time-scales in order to determine accurately the various lamp-post model parameters; something which is beyond the sensitivity limit of current X-ray observatories.\\par Both of these observational problems will be solved by future X-ray observatories. Athena, due for launch in 2028, will be placed at the L2 point and so will allow continuous uninterrupted observations for as long as it is needed, in order to sample completely the low frequency part of the time-lag spectra of the high BH mass AGN. Athena will also have almost an order of magnitude greater sensitivity than \\textit{XMM}-\\textit{Newton}, allowing detailed measurements of the time-lags in the high frequency domain. Moreover, if approved, LOFT, although not allowing long continuous observations, will improve sensitivity at high frequencies due to its superior effective area.\\par A major result from our analysis is that the height of the X-ray source must be very close to the BH, around 3.7 \\rg, in all AGN. In the lamp-post geometry layout this X-ray source should somehow correspond to the centroid of a hemispherical X-ray corona lying above the accretion disc (a similar X-ray source exists below the disc) where the bulk of the hard X-rays are produced. The derived compactness of the X-ray region and its proximity to the central BH is in accordance with recent X-ray microlensing studies \\citep{chartas09,chartas12}, who also find values for the X-ray source size as small as 6--10 \\rg. The lack of correlation between the height and either the mass and/or the BH spin indicates that the dimensions of the X-ray region must by tuned by some other physical parameters, perhaps the (currently unknown) physical mechanism which generates this X-ray emission, which should be the same in all objects. Note that if there was any height-luminosity dependence, it would have been wiped out since we are using luminosity-averaged time-lag spectra. In a future work we will address this issue by modelling the time-lag spectra during high and low source states.\\par Another major result from our work is that the BH spin parameter, $\\alpha$, does {\\it not} appear to be the same in all objects. We find that five sources in our sample, namely Mrk\\;766, MCG--6-30-15, IRAS\\;13224-3809, ESO\\;113-G010 and NGC\\;5548, possibly host a rotating BH, since the best-fitting $\\alpha$ for these objects differs significantly from zero. On the other hand, the best-fitting spin parameters differ significantly from unity in the case of NGC\\;4051, Ark\\;564, 1H\\;0707-495 and NGC\\;3516 suggesting that they may host a non-rotating BH. \\citet{reynolds13} considered a sample of 20 AGN and, by compiling X-ray spectral fitting from the literature, found high spin parameter values ($\\alpha>0.8$) for $M=\\left(2\\times 10^6-3\\times10^7\\right)$ \\ms, with a hint of lower values ($\\alpha\\simeq0.5$) for $M<2\\times 10^6$ and $M>5\\times10^7$ \\ms, favouring a prolonged ordered (coherent) accretion scenario that spins up the BH to high values. Our results are in agreement with theirs, although we do not observe a correlation between BH spin and mass. This may be due to the rather small size of our sample. On the other hand, our results i.e.\\ high and low spin parameters and the lack of correlation between $\\alpha$ and $M$, favour a galaxy mergers scenario, in which the BH growth occurs in gas-poor galaxies, resulting to a distribution of spins that has little dependence on BH mass \\citep{volonteri13}.\\par Our work shows that derivation of the geometry of the X-ray source and surrounding reprocessing material (here modelled as an accretion disc) in the vicinity of BHs is possible from studies of the Fourier-resolved time-lags, between soft and hard X-ray bands. In future papers, we will study, via GR ray tracing, X-ray photons of all energies as they are reflected and reprocessed by the accretion disc. Not only will we then be able to properly trace the behaviour of soft photons, after they are emitted from the disc, but we will also be able to determine the contribution of reflected photons to the observed 1.5--4 keV hard X-ray band. At the same time, we will model separately the high and low X-ray flux states of all AGN, in order to trace possible luminosity dependencies with respect to the X-ray source height. With future X-ray observatories such as Athena and LOFT it will be possible to improve significantly on our present time-lag spectral estimates by extending them to a much wider mass range, decreasing at the same time their uncertainties."
    },
    "1402/1402.0407_arXiv.txt": {
        "abstract": "The hierarchical triple system PSR~J0337+1715 offers an unprecedented laboratory to study secular evolution of interacting systems and to explore the complicated mass-transfer history that forms millisecond pulsars and helium-core white dwarfs.  The latter in particular, however, requires knowledge of the properties of the individual components of the system.  Here we present precise optical spectroscopy of the inner companion in the PSR~J0337+1715 system.  We confirm it as a hot, low-gravity DA white dwarf with $\\Teff=15,800\\pm100\\,$K and $\\logg=5.82\\pm0.05$. We also measure an inner mass ratio of $0.1364\\pm0.0015$, entirely consistent with that inferred from pulsar timing, and a systemic radial velocity of $29.7\\pm0.3\\,\\kms$.  Combined with the mass ($0.19751\\,M_\\odot$) determined from pulsar timing, our measurement of the surface gravity implies a radius of $0.091\\pm0.005\\,R_\\odot$; combined further with the effective temperature and extinction, the photometry implies a distance of $1300\\pm80\\,$pc.  The high temperature of the companion is somewhat puzzling: with current models, it likely requires a recent period of unstable hydrogen burning, and suggests a surprisingly short lifetime for objects at this phase in their evolution.  We discuss the implications of these measurements in the context of understanding the PSR~J0337+1715 system, as well as of low-mass white dwarfs in general. ",
        "introduction": "\\label{sec:intro}   White dwarfs (WDs) are among the best-understood stars, enabling their use as astrophysical tools in investigations of, e.g., the ages and masses of astrophysical systems \\citep[e.g.,][]{acigb10}.  However, the lowest-mass WDs with He cores --- Extremely-Low Mass (ELM) WDs with masses $<0.2\\,M_\\odot$ \\citep[][and references therein]{kbap+12,bkap+13} --- still defy complete understanding, with few reliable independent measurements of masses, sizes and ages known \\citep[e.g.,][]{vkbk96,bsp+05}.  Yet, these properties are  important for understanding the evolution of ELM WDs and the binaries they are found in \\citep*{il93,mdd95}. For instance, while one would naively expect low-mass WDs to cool quickly, given their relatively large size and small heat capacity, some ELM WDs can remain bright and hot \\citep[cf.][]{llfn95} because they have outer hydrogen layers sufficiently thick for nuclear fusion to continue -- stably or otherwise -- for several Gyr (\\citealt{asvdhp96,sarb02,pach07}; \\citealt*{ambc13}).  Improving our understanding of ELM WD cooling would aid in evolutionary models for, e.g., millisecond pulsars and the later stages of mass transfer \\citep[e.g.,][]{tlk12,avkk+12,afw+13,kbvk+13}. Similarly, improved masses and radii would aid in determining the final fates of double-WD binaries (\\citealt{db03}; \\citealt*{mns04}; \\citealt{davb+06}; \\citealt*{kbs12}): R~CrB stars, AM~CVn binaries, or even SNe Ia \\citep{it84,webbink84}.     \\object[PSR J0337+1715]{PSR~J0337+1715} (hereafter \\psr; \\citealt{rsa+14}) was discovered in the 350\\,MHz Green Bank Telescope Driftscan survey \\citep{blr+13,lbr+13}, and initial timing observations found a $2.7\\,$ms spin period, a $1.6\\,$d orbital period, and a likely companion mass of 0.1--0.2\\,$M_\\odot$, all of which are consistent with expectations for a fully-recycled pulsar with a low-mass He WD companion \\citep{vkbjj05,tlk12}.  However, further deviations to the observed spin period soon became apparent, suggesting the presence of an additional body in the system.  This was confirmed with an intensive timing campaign, finding an outer orbital period of 327\\,d \\citep{rsa+14}.  By comparing the pulsar's pulse arrival times with numerical integrations of possible orbits, the masses, inclinations, and orbital parameters of the system could be determined precisely \\citep[][Archibald et al.\\ 2014, in prep.]{rsa+14}, with a pulsar mass of $1.4378\\pm0.0013\\,M_\\odot$, an inner companion mass of $0.19751\\pm0.0015\\,M_\\odot$, and an outer companion mass of $0.4101\\pm0.0003\\,M_\\odot$; both inner and outer orbits are inclined by about $39\\degr$ to the plane of the sky.   Before the nature of the system was fully determined, we identified an unusually blue object coincident with it, which, based on initial spectroscopy and photometry, we identified as the inner companion, almost certainly a hot low-mass WD \\citep{rsa+14}.  The brightness of the system, its proximity, and the detailed constraints offered by pulsar timing make the system a fantastic laboratory to explore the atmosphere, structure, and evolution of ELM WDs.  Here, we present an intensive spectroscopic campaign aimed to do so, which also serves as a valuable cross-check of the  pulsar timing results. ",
        "conclusions": "Our measurements provide the velocities and atmospheric parameters of the inner white dwarf in the \\psr\\ system.  The velocities serve to confirm the more precise ones inferred from pulsar timing.  In addition, they show that the systemic velocity is low.  This is not unexpected, since any kick imparted to the system in the supernova explosion that formed the neutron star must have been small for the triple to survive \\citep{tvdh14}.  One thus expects the proper motion to be similarly small.  Our measurement of the surface gravity, combined with the precise mass from timing, implies a radius of $0.091\\pm0.005\\,R_\\odot$.  Combining this in turn with the effective temperature, extinction, and photometry, one infers a distance of $1300\\pm80\\,$pc \\citep{rsa+14}. With an accurate parallax from very-long baseline interferometry (measurements are in progress), this can be used to infer the surface gravity and thus test the model atmospheres in an otherwise poorly constrained regime.  The mass and radius of the inner white dwarf are consistent with those expected for a young, low-mass helium-core white dwarf, similar to the white dwarf companions found in other binaries (Section~\\ref{sec:intro}).  Compared to low-mass white dwarfs around pulsars,\\footnote{One cannot easily compare  low-mass white dwarfs with white-dwarf or A-star companions, since for those systems there are strong observational biases to find hotter, more luminous and larger white dwarfs.} however, the source stands out for being hotter than most.  This must be intrinsic as possible contributions from pulsar irradiation and tidal heating are negligible.  The high temperature is surprising as it would suggest the system is in a short-lived state and hence that similar systems are common -- which, empirically, they are not.  This suggestion arises because in current evolutionary models of helium white dwarfs, temperatures in excess of $\\sim\\!12,000\\,$K are only achieved in models with unstable shell flashes (e.g., \\citealt{dsbh98}; \\citealt*{asb01}).  In those shell flashes, however, most of the thick hydrogen layer is lost, and hence the white dwarf will cool relatively quickly.  Furthermore, while the flashing state may last $\\sim\\!200\\,$Myr \\citep{ambc13}, the typical cooling timescales at these temperatures are short, a few $10\\,$Myr for each flash, and hence the total time spent at high temperatures is often $<100\\,$Myr, depending on mass.  These timescales are still far longer than the expected sedimentation timescale for helium ($\\sim 10^3\\,$yr) following mixing during a shell flash, consistent with the lack of any \\ion{He}{1} in the spectrum of the inner WD (cf.\\ \\citealt{kbvk+13}).  Based on inspection, we can roughly limit its abundance to $10^{-2.5}$H (by number), which would not change our inferred \\logg\\ by more than our quoted uncertainty.  If the inner white dwarf is really only a few 100\\,Myr old, it almost certainly formed last, as also expected from simple models \\citep{rsa+14,tvdh14}, although we cannot strictly exclude the opposite: at the upper limit to the temperature of the outer WD of 20,000\\,K \\citep{rsa+14}, the cooling time would be 30--100\\,Myr.  Of course, it could be a coincidence that we found such a hot white-dwarf companion.  However, a similarly hot companion was found for \\object[PSR J1816+4510]{PSR~J1816+4510} \\citep{kbvk+13}.  Since typical millisecond pulsars remain visible for a Hubble time, and since we know $\\sim\\!50$ pulsars with low-mass white dwarf companions, this suggests that white dwarfs can stay hot for $\\sim\\!500\\,$Myr, substantially longer than expected in current theoretical models.  The discrepancy is made worse by the fact that white dwarfs with mass below $\\lesssim\\!0.18\\,M_\\odot$ are not expected to get this hot at all (they should not have flashes).  Observationally, however, lifetimes of a few hundred Myr seem also indicated by the prevalence of hot, low-mass white dwarfs around A stars (which have ages of $\\lesssim\\!1\\,$Gyr)."
    },
    "1402/1402.7227_arXiv.txt": {
        "abstract": "Dust in debris discs is constantly replenished by collisions between larger objects. In this paper, we investigate a method to detect these collisions. We generate models based on recent results on the Fomalhaut debris disc, where we simulate a background star transiting behind the disc, due to the proper motion of Fomalhaut. By simulating the expanding dust clouds caused by the collisions in the debris disc, we investigate whether it is possible to observe changes in the brightness of the background star. We conclude that in the case of the Fomalhaut debris disc, changes in the optical depth can be observed, with values of the optical depth ranging from $10^{-0.5}$ for the densest dust clouds to $10^{-8}$ for the most diffuse clouds with respect to the background optical depth of $\\sim1.2\\times10^{-3}$. ",
        "introduction": "Debris discs are circumstellar belts of dust and debris around stars. These discs are analogous to the Kuiper belt and the asteroid belt in our own Solar system. They provide a stepping stone in the study of planet formation, because the evolution of a star's debris disc is indicative of the evolution of its planetesimal belts \\citep{/disks/strw1/zeegers/major_project/papers/Wyatt08}.  Debris discs can be found around both young and more evolved stars. For the youngest stars, the dust in the disc can be considered a remnant of the protoplanetary disc. In these young debris discs, gas might still be present, which is not the case for debris discs around more evolved (main sequence) stars. For older stars, it is more likely that debris discs indicate the place where a planet has failed to form. This can either be because the formation time-scale was too long, like in the Solar system's Kuiper Belt, or because the debris disc was stirred up by gravitational interactions of other planets before a planet could form, which probably happened in the Solar system's Asteroid belt \\citep{/disks/strw1/zeegers/major_project/papers/Wyatt08}. The dust in debris discs is thought to be constantly replenished by collisions between the planetesimals. These planetesimals start to grow in the disc during the protoplanetary disc phase. Models indicate that when the planetesimals begin to reach the size of $\\sim2000$ km in diameter, the process of growth reverses and the disc begins to erode. The dynamical perturbations of these large objects (with diameters $>$ 2000 km) stir up the disc and start a cascade of collisions ~\\citep{/disks/strw1/zeegers/major_project/papers/Kenyon05}. However, it is not clear from these models how the dust is replenished over a time-scale of more than 100 Myr \\citep{/disks/strw1/zeegers/major_project/papers/Wyatt08}.   The first debris discs were discovered with the Infrared Astronomical Satellite (\\emph{IRAS}; \\citep{/disks/strw1/zeegers/major_project/papers/Neugebauer84}), which measured the excess emission in the infrared caused by dust in the debris disc. The dust is heated by the central star and therefore re-emits thermal radiation, which causes the observed spectrum of the system to deviate from that of a stellar black body radiation curve. The debris disc around Vega was the first debris disc discovered in this way \\citep{/disks/strw1/zeegers/major_project/papers/Aumann84} and after that more than 100 discs have been subsequently discovered. Observations from recent surveys indicate that at least 15 per cent of FGK stars and 32 per cent of A stars have a detectable amount of circumstellar debris (\\citet{2014MNRAS.437.3288B}, \\citet{/disks/strw1/zeegers/major_project/papers/Bryden06}, \\citet{2007ApJ...658.1312M}, \\citet{2008ApJ...677..630H},\\citet{2009MNRAS.397..757G} and \\citet{2006ApJ...653..675S}).   Until improved coronagraph techniques became available, the only ground-based resolved example of a debris disc observed in scattered light at an optical wavelength of $0.89\\,\\mu$m was the debris disc of Beta Pictoris \\citep{/disks/strw1/zeegers/major_project/papers/Smith84}. However, during the past decade many resolved debris discs have been observed at optical and near-infrared wavelengths. Debris discs have been observed in scattered light using the Advanced Camera for Surveys (ACS; \\citep{2004IAUS..221..449C}) as well as the Space Telescope Imaging Spectrograph (STIS) and in the near-infrared ($1.1\\,\\mu$m) using the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) combined with the usage of coronagraphs on the \\emph{Hubble Space Telescope (HST)}. Examples of debris disc observed with these instruments are: the debris disc of HD 202628 observed with STIS (see Figure~\\ref{fig:HD 202628} and \\citet{/disks/strw1/zeegers/major_project/papers/Krist12}), the debris disc of AU Microscopii \\citep{/disks/strw1/zeegers/major_project/papers/Krist05} with the ACS and the NICMOS image of the debris disc HR 4769A \\citep{1999ApJ...513L.127S}. Debris discs are also observed at infrared and (sub)millimetre wavelengths where the dust emits the reprocessed stellar light as thermal radiation, for example the debris disc of Epsilon Eri observed at $850\\,\\mu$m with the Submillimetre Common-User Bolometer Array (SCUBA) at the James Clerk Maxwell Telescope \\citep{1998ApJ...506L.133G} and Beta Pictoris observed with Herschel Photodetector Array Camera \\& Spectrometer (PACS) and Spectral and Photometric Imaging Receiver (SPIRE; \\citep{2010A&A...518L.133V}). These direct observations of debris discs show a wide variety of disc morphology. Some of these discs have narrow dust rings, while other discs are more widespread. Systems can have multiple and even warped discs. Models show that the shape of debris discs can be caused by shepherding planets around these rings (\\citet{/disks/strw1/zeegers/major_project/papers/Deller05}; \\citet{/disks/strw1/zeegers/major_project/papers/Quillen07}). The first hints that this might be the case are given by observations of $\\beta$ Pictoris (\\citet{/disks/strw1/zeegers/major_project/papers/Lagrange10}; \\citet{/disks/strw1/zeegers/major_project/papers/Quanz10}) and HD 100456 \\citep{/disks/strw1/zeegers/major_project/papers/Quanz13}.  One such a star with a debris disc is the nearby [$7.668 \\pm 0.03$ pc \\citep{/disks/strw1/zeegers/major_project/papers/Perryman97}] A star Fomalhaut. The most prominent feature of the disc is the main dust ring at a radius of 140 au from the star \\citep{/disks/strw1/zeegers/major_project/papers/Kalas05}, which is $\\sim25$ au wide. This debris disc has been imaged by the \\emph{HST} in 2005 \\citep{/disks/strw1/zeegers/major_project/papers/Kalas05} and more recently by the Herschel Space Telescope \\citep{/disks/strw1/zeegers/major_project/papers/Acke12_updated} and the Atacama Large Milimeter Array (ALMA) \\citep{/disks/strw1/zeegers/major_project/papers/Boley12}. The dust around the Fomalhaut debris disc has been observed in both reflected optical light and at $10 -100 \\mu \\mathrm{m}$ wavelength, where the thermal radiation of the dust in the disc is observed radiation \\citep{/disks/strw1/zeegers/major_project/papers/Holland98}.  The ring has a mass-loss rate of $2\\times10^{21} \\,\\mathrm{g\\,yr}^{-1}$ \\citep{/disks/strw1/zeegers/major_project/papers/Acke12_updated}, which can be compared to the loss of the total mass of the rings of Saturn per year. This huge amount of mass suggests a high collision rate. \\citet{/disks/strw1/zeegers/major_project/papers/Wyatt02} investigated the possibility of large dust clumps in the Fomalhaut debris disc due to collisions between large planetesimals ($>1400 \\,\\mathrm{km}$) in order to explain a residual arc of 450$\\,\\mu$m emissions approximately 100 au from the star. These collisions would make an observational detectable clumpy morphology. Such a dust clump may be detected in the debris disc of Beta Pictoris. \\citet{1995A&A...299..557L} conclude in their paper that the brightness variation in this star on a time-scale from 1979 until 1982 can be attributed to either occultation of the star by a clumpy dust cloud or a planet. However, recent observations of the Fomalhaut debris disc in thermal emission show a smooth structure to the debris, which hints at a high dust replenishment rate by numerous collisions \\citep{/disks/strw1/zeegers/major_project/papers/Acke12_updated}. The colliding planetesimals will have diameters smaller than 100 km (\\citet{/disks/strw1/zeegers/major_project/papers/Wyatt02}; \\citet{/disks/strw1/zeegers/major_project/papers/Greaves04}; \\citet{/disks/strw1/zeegers/major_project/papers/Quillen07}) and therefore the dust clouds will be difficult to detect either in reflected optical light or at longer wavelengths.   Current observations of debris discs show us the distribution of small dust particles, with radii from $10^{-5}$ m to 0.2 m \\citep{/disks/strw1/zeegers/major_project/papers/Wyatt02}. We are not able to directly observe planetesimals or large boulders. This means that we do not have an observational confirmation of the distribution of these larger parent particles. Observing the debris resulting from collisions would make it possible to put constraints on the particle-size distribution in debris discs. In this paper, we explore a technique that would make it possible to indirectly observe collisions between large planetesimals. When a background object, like a star, passes behind a debris disc, dust generating collisions will cause the star to dim slightly. The change in brightness depends on the optical depth of the dust clouds, which in turn depends on the amount of debris created in the collision between two planetesimals.  The outline of the paper is as follows. Section 2 explains the method we use to observe collisions in debris discs in more detail and gives a general introduction to the Fomalhaut debris disc. In Section 3, we explain the theoretical background of the collision model used. Section 4 explains the difference between three models of collisions in the debris disc of Fomalhaut. Section 5 shows differences in the observations for systems with different inclinations. In Section 6 we show other debris discs with background objects passing behind the disc, and we conclude in Section 7 with a summary and a discussion of our results. ",
        "conclusions": "In this paper, we simulate a background star transiting behind the Fomalhaut debris disc. This background star gives us an opportunity to observe the distribution of planetesimals as a function of position in the debris disc and to test the collisional cascade model, which determines the size distribution of the planetesimals according to theoretical models. We modeled these collisions using three different models that differ in initial size distribution of the planetismals and in size distribution of the debris particles. We find that if all the colliding particles are pulverized 1-km-sized planetesimals, the value of the variations of the optical depth due to the dust clumps caused by the collisions ranges between $10^{-2.5}-10^{-1.5}$, establishing a strong upper limit to the experiment. To make our collision model more realistic, we introduced a size distribution based on the collisional cascade. The result of the second scenario shows that the variations of the optical depth decreases to values of $10^{-6.5}-10^{-4.5}$. The third scenario includes a largest remnant particle. Half the mass of the debris particles remains locked up in the largest remnant. Again the value of $\\tau$ drops, but also the number of detections drops dramatically. These low mass debris clouds will be impossible to observe.  We should be able to measure changes in the flux from a background star caused by catastrophic collisions between large boulders of debris in the disc with diameters of 1 km or larger. It is not possible to have the \\emph{HST} continuously pointing at Fomalhaut, but this is not necessary. From Fig.~\\ref{fig:frecuency_observing} we conclude that there is a 10 per cent chance that the observed optical depth changes significantly after 100-150 d and from Fig.~\\ref{fig:frecuency_observing_100m} we conclude that the optical depth changes significantly after 30-50 d. It would therefore suffice to observe the debris disc once every 1-2 months.   The values of the optical depth resulting the simulations of our second and third scenario are below the detection limits of current telescopes. However, we emphasize that these models are based on the assumption that the size distribution of the debris resulting from the collision follows a collisional cascade. The largest particles in the cascade are 10 meter boulders (scenarios 1 and 2) or the second largest remnant particle (scenario 3) and the smallest particles are 8 $\\mu$m particles, just above the blow-out size. The total debris mass is equal to the mass-loss rate, which was determined for Fomalhaut by \\citet{/disks/strw1/zeegers/major_project/papers/Acke12_updated}. They show that this mass-loss rate can be compared to 1000 collisions between 1-km-sized planetesimals, where these planetesimals pulverize to small dust particles instead of letting the size distribution of the debris follow a collisional cascade. Using this particle distribution we underestimate the amount of dust generated in the debris disc, but in this way we do not overestimate the number of collisions between large 1-km-sized planetesimals. This does mean that there are very likely more collisions in the disc than we assume in this paper, but it is not clear what quantity of dust is produced in cratering collisions and what is produced in catastrophic collisions. We emphasize that we used the worst case scenario in this paper and even in these cases it is not impossible to detect a change in the optical depth. Future simulations will therefore have to involve a more realistic collision rate, where the mass-loss rate is built up by the amount of blow-out particles that arise from a collision. It is important to make some strong assumptions on the contribution of the amount of dust resulting from cratering and from catastrophic collisions between particles smaller than 100 m.   It is not impossible that we will measure optical depths caused by more than one collision, due to the overlap of at least two dust clouds resulting from collisions. Our simulations can distinguish between two or more dust clouds, but in reality this will be harder. The differences in optical depth in the case of overlap is often very large. In the cases we encountered, it involved a far expanded dust cloud with a very low value of $\\tau$ of the order of $10^{-7}$ and a much smaller and denser dust cloud. In reality, this would mean that if we observe a change in the optical depth, most of the change is due to the more compact dust cloud and the signal of the other cloud will be too weak to detect.  Besides the Fomalhaut debris disc, we investigated whether there are more debris discs with a background object that will transit in the next decade. We concluded that there are three other debris discs with a background object transiting behind the disc: HD 107146, HD 202628 and AU Microscopii. The star HD 107146 has a galaxy that will move behind their discs and AU Microscopii has both a star and galaxy transiting behind the disc, but in the case of the star HD 202628 it is not certain whether a galaxy or a star will pass behind the disc. We are looking for edge-on systems in particular, because for these systems we can look through two sides of the debris ring and $\\tau$ increases.  Using a background star to indirectly observe collisions in debris discs is a powerful method to investigate the size distribution of debris in debris discs. Though we use a very simple model, it is now possible to put some boundaries to the detectability of the dust clouds resulting from the collisions between planetesimals. When observations prove that there is no change in the brightness of the background star, we can conclude that most of the dust is created in cratering collisions or in catastrophic collisions that have a low impact energy whereby at most half of the mass of the planetesimal remains intact.  Future work includes the role of gas absorption in the line of sight and modeling of other candidate systems."
    },
    "1402/1402.5963_arXiv.txt": {
        "abstract": "Magnetic fields play an important role in many astrophysical systems and a detailed understanding of their impact on the gas dynamics requires robust numerical simulations. Here we present a new method to evolve the ideal magnetohydrodynamic (MHD) equations on unstructured static and moving meshes that preserves the magnetic field divergence-free constraint to machine precision. The method overcomes the major problems of using a cleaning scheme on the magnetic fields instead, which is non-conservative, not fully Galilean invariant, does not eliminate divergence errors completely, and may produce incorrect jumps across shocks. Our new method is a generalization of the constrained transport (CT) algorithm used to enforce the $\\nabla\\cdot \\mathbf{B}=0$ condition on fixed Cartesian grids. Preserving $\\nabla\\cdot \\mathbf{B}=0$ at the discretized level is necessary to maintain the orthogonality between the Lorentz force and $\\mathbf{B}$. The possibility of performing CT on a moving mesh provides several advantages over static mesh methods due to the quasi-Lagrangian nature of the former (i.e., the mesh generating points move with the flow), such as making the simulation automatically adaptive and significantly reducing advection errors. Our method preserves magnetic fields and fluid quantities in pure advection exactly. ",
        "introduction": "\\label{sec:intro}  The equations of fluid dynamics and magnetohydrodynamics (MHD) may be evolved using a variety of numerical approaches. One can choose from a number of discretizations that are formally intended to yield a solution that is accurate to some order $n$. However, a system of partial differential equations for physical flows also gives rise to conserved quantities due to the symmetries of the problem and these quantities are not necessarily preserved with some algorithms. Such conserved quantities for the MHD equations include mass, momentum, angular momentum, energy, and the solenoidal nature of the magnetic field. It then becomes a concern to design techniques in a clever way such that the discrete representations of these conserved quantities are in fact preserved to machine precision. Failure to do this may result in instabilities and unphysical solutions; that is, the system's quantity that should be conserved may chaotically evolve away from its initial value. One well-known example of such a phenomenon occurs when evolving an $N$-body system with a non-symplectic time integrator, which does not conserve the total energy and hence may lead to energy drifts and decaying orbits.  In this work, we propose a scheme for evolving the MHD equations on a moving mesh such that conservation of mass, momentum, and energy, and, importantly, the solenoidal (divergence-free) nature of the magnetic field is preserved. On static Cartesian meshes, the constrained transport (CT) algorithm achieves this goal. It uses a finite volume formalism to evolve the density, momentum, and energy, and exploits Stokes' theorem and uses a face-averaged representation of the magnetic fields (called the `staggered-mesh' approach) to enforce $\\nabla\\cdot \\mathbf{B}=0$ \\citep{1988ApJ...332..659E}. The CT algorithm has been described in the literature as being quite difficult (if not impossible) to extend to an unstructured mesh \\citep{2011ApJS..197...15D,2011MNRAS.418.1392P,2013MNRAS.432..176P} and leading to the development of alternate divergence-cleaning schemes to keep $\\nabla\\cdot \\mathbf{B}$ small but non-zero on moving meshes based on the Powell \\citep{1999JCoPh.154..284P} and Dedner \\citep{2002JCoPh.175..645D} cleaning methods. Moving mesh methods for MHD with cleaning schemes have been described by \\cite{2011ApJS..197...15D,2011MNRAS.418.1392P,2012ApJ...758..103G,2013MNRAS.432..176P}. Additionally, there have been recent advances in designing robust mesh-less based methods for MHD \\citep{2011MNRAS.414..129G,2013MNRAS.436.2810T}, which still, however, require cleaning schemes. These schemes alter the MHD equations and may add source terms in an attempt to control the divergence errors, which has the unwanted side-effect of making the schemes non-conservative. Additionally, such an approach loses the moving mesh's Galilean invariant properties (e.g. solving pure advection problems is diffusive). One approach to improve these methods has been to develop a {\\it locally} divergence-free formulation based on the discontinuous Galerkin method in the work presented in \\cite{2014MNRAS.437..397M}, where local divergence-free basis functions are used to represent the solution on each cell.  In this paper we solve the outstanding problem of designing a {\\it globally} divergence-free, conservative, finite volume based algorithm on an unstructured moving mesh.  Coupling a CT algorithm with a moving mesh technique is clearly highly desirable. The moving mesh algorithm for solving the inviscid Euler equations was developed recently by \\cite{2010MNRAS.401..791S} and implemented in the code {\\sc arepo}. The moving mesh algorithm largely eliminates a number of well-know weaknesses of both static Cartesian/adaptive mesh refinement (AMR) approaches and pseudo-Lagrangian smoothed-particle hydrodynamics (SPH). The mesh-generating points in the algorithm can be set to move with the fluid flow, making the scheme quasi-Lagrangian. The method has automatic adaptivity, is Galilean-invariant, and significantly reduces advection errors. Contact discontinuities resulting from shocks are preserved to significantly greater accuracy. In addition, conditions on the size of the time-step are less severe than on a static mesh in the cases of flows with high-Mach number bulk motions, since fluxes across cells are always calculated in the rest-frame of the faces. A method for solving the MHD equations on a moving mesh will automatically gain these benefits of the Lagrangian nature of the scheme and, ideally, should not suffer any significant disadvantage compared to a static mesh approach. Since methods exist to preserve the divergence-free condition on structured fixed grids, it is important to extend this to unstructured meshes so that if differences are observed in MHD simulations run on static and moving meshes, then they can be attributed to the advantages of the moving mesh approach and not as an artefact produced by divergence errors.  The CT method proposed here can also be applied to static arbitrarily unstructured grids, where CT algorithms are lacking. A number of simulations use non-Cartesian static grids, dictated by the geometry/symmetry of the problem, as a way of achieving greater accuracy and by using fewer grid cells. For example, the MHD solver in \\cite{2013ApJS..205...19F} was developed for a hexagonal spherical geodesic grid as a way to improve the simulation of astrophysical flows of partially ionized plasmas around a central compact object. However, such methods use cleaning schemes instead of CT to handle the magnetic field, and would benefit greatly from the CT method described in what follows.  A robust code for solving the MHD equations as accurately as possible would have many astrophysical applications because magnetic fields are ubiquitous in the Universe. For example, the magnetorotational instability in accretion discs \\citep{1998RvMP...70....1B} generates turbulence and mediates angular momentum transfer. Discs around supermassive black holes may be levitated by magnetic pressure \\citep{2012ApJ...758..103G}. Magnetic fields are key in the production of relativistic jets and outflows from compact sources \\citep{1977MNRAS.179..433B,1982MNRAS.199..883B}. Magnetic fields under a vertical gravitational field as in the disc of a galaxy leads to the Parker instability, which is thought to play an important role in the evolution of the interstellar medium \\citep{1966ApJ...145..811P,1967ApJ...149..535P}. Interstellar turbulence and star formation are linked with magnetized plasma processes \\citep{1995ApJ...438..763G}. Cosmic magnetic fields are likely to have a primordial origin which leads to imprints on the temperature and polarization anisotropies of the cosmic microwave background radiation \\citep{2001PhR...348..163G}. These seed fields are amplified during the formation and evolution of galaxies. Radio lobes of galaxies may also play an important role in enhancing the magnetic fields in the intergalactic medium \\citep{2001ApJ...560L.115G}. Magnetic fields may be responsible for suppressing strong isotropic turbulence and conduction in clusters and make it possible for stable, $100$ kpc-scale high-density arms to exist in the cluster environment \\citep{2013Sci...341.1365S}. The strength and topology of magnetic fields are responsible for determining the propagation of cosmic rays in galaxies \\citep{1998ApJ...509..212S}. Magnetic fields are also present in a wide variety of stars and play a significant role in their evolution \\citep{2009ARA&A..47..333D}.  Many of these problems are suitable for study by a moving mesh approach and up until now a divergence-free MHD solver has been lacking. This paper lays the framework for implementing a divergence-preserving algorithm for robust evolution of the MHD equations in moving mesh codes such as {\\sc arepo} \\citep{2010MNRAS.401..791S} and {\\sc tess} \\citep{2011ApJS..197...15D}.  In Section~\\ref{sec:methods} we describe the details of the numerical method. In Section~\\ref{sec:tests} we show the results of numerical tests (with comparisons to fixed grid CT and the Powell cleaning scheme on a moving mesh) demonstrating that the method works well and has several advantages over the other two techniques. In Section~\\ref{sec:disc} we discuss variations of our method and future directions. In Section~\\ref{sec:conclusion} we provide concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a new, stable, accurate, and robust CT scheme for solving the MHD equations on a moving unstructured mesh which is conservative and preserves the divergence of the magnetic field to zero at the level of machine precision. Such a CT scheme for a moving mesh has been claimed to be difficult to construct, maybe even impossible, by other authors \\citep{2011ApJS..197...15D,2011MNRAS.418.1392P,2013MNRAS.432..176P}, but we have demonstrated that the scheme can be achieved. The CT method significantly improves the other current methods used to evolve the MHD equations on unstructured meshes and moving meshes, which are not necessarily conservative and can show large divergence errors and incorrect shock jump conditions. The new numerical method also has significant advantages over CT approaches on a fixed grid. Namely, our method is a quasi-Lagrangian scheme and greatly reduces advection errors. In pure advection flows, it preserves the solution at the level of machine-precision, which is a uniquely powerful feature of the method. Due to the moving mesh formulation, the method is also automatically adaptive in its resolution. Additionally, Galilean boosts/large velocities in the flow affect the truncation errors to a significantly smaller extent than in fixed grid codes due to the Galilean invariant properties of the moving mesh code. It is vital to construct a moving mesh code that preserves $\\nabla\\cdot \\mathbf{B}=0$ to machine precision, otherwise if differences between moving and fixed grid MHD codes are observed it becomes difficult to tell whether it is due to the advantages of moving mesh codes (such as better treatment of advection) or to the non-zero divergence errors that exist in the solution.  The new method offers new exciting possibilities of simulating MHD flows more accurately than before. One particular application it is well-suited for is the study of supersonic MHD turbulence. Here the flow has high Mach number velocities so fixed grid codes experience very restrictive timestep criteria, which can be lessened on a moving mesh. Additionally, advection errors would be reduced to a significant degree in moving mesh CT simulations. An analytic theory for supersonic MHD turbulence is virtually non-existent and numerical simulations are considered the best framework for studying this problem, necessitating numerical methods that minimize any unphysical effects of numerical errors on the solution. The moving mesh CT approach is a very viable candidate for this task.  The moving mesh CT method is also well-suited for studying astrophysical problems where advection is important. Most simulations, from accretion disc simulations to cosmological simulations, have regions that are only marginally resolved due to limitations in computational resources. As we see from our low-resolution simulations of advection, the solution in static CT methods and moving mesh methods with divergence cleaning schemes can have large diffusion errors which can be reduced only with higher resolution simulations. The moving CT approach, however, is able to accurately advect the solution to machine precision, which is one of the main benefits of the method."
    },
    "1402/1402.2896_arXiv.txt": {
        "abstract": "We analysed  archival data of \\textit{Chandra} pre-explosion observations of the position of SN2014J in M82. No X-ray source at this position was detected in the data, and we calculated upper limits on the luminosities of the progenitor. These upper limits allow us to firmly rule out an unobscured supersoft X-ray source progenitor with a photospheric radius comparable to the radius of white dwarf near the Chandrasekhar mass ($\\sim$1.38 M$_{\\odot}$) and mass accretion rate in the interval where stable nuclear burning can occur. However, due to a relatively large hydrogen column density implied by optical observations of the supernova, we cannot exclude a supersoft source with lower temperatures, $kT \\lesssim 80$ eV.  We find that the supernova is located in the center of a large structure of soft diffuse emission, about 200 pc across. The mass, $\\sim 3\\cdot 10^4$ M$_\\odot$ and short cooling time of the gas, $\\tau_{\\rm cool}\\sim 8$ Myrs, suggest that it is a supernova-inflated super-bubble, associated with the region of recent star formation. If SN2014J is indeed located inside the bubble, it likely belongs to the prompt population of type Ia supernovae, with a delay time as short as $\\sim 50$ Myrs. Finally, we analysed the one existing post-supernova \\textit{Chandra} observation and placed upper limit of $\\sim (1-2)\\cdot 10^{37}$ erg/s on the X-ray luminosity of the supernova itself. ",
        "introduction": "\\label{Sect:Introduction} Type Ia supernovae are highly luminous stellar explosions, visible over cosmological distances. The chemically enriched ejecta from supernova explosions are important to galactic evolution, and may also play a role in the enrichment of the intergalactic medium \\citep{Hillebrandt.et.al.2013}. The correlation between peak brightness and exponential fall-off time of the light-curves of type Ia supernovae (\\citealt{Philips.1993}) means that they can be used as standardisable cosmological candles to measure the expansion history of the Universe \\citep{Riess.et.al.1998,Perlmutter.et.al.1999}. However, the progenitor systems giving rise to these supernovae remain poorly understood, despite decades of studies (see e.g. \\citealt{Maoz.et.al.2013}). The consensus is that the exploding objects are carbon-oxygen white dwarfs near the Chandrasekhar mass ($\\sim$1.38 M$_{\\odot}$) that undergo thermonuclear runaways when the density and pressure in their cores reach values where carbon and oxygen will be processed unstably into radioactive iron-group elements. The resulting explosions completely destroy the white dwarfs, and the subsequent decay of mainly radioactive $^{56}$Ni in the ejecta powers exponential light-curves which, for a period of weeks to months may outshine the integrated luminosity of the host galaxies. However, it remains unclear how a carbon-oxygen white dwarf may reach the Chandraskhar mass. Carbon-oxygen white dwarfs are usually formed at masses significantly lower than that required for thermonuclear runaway, and since there is no known way in which a single white dwarf may grow in mass a binary origin of type Ia supernova progenitors is usually assumed. Two general classes of progenitor scenarios are commonly considered: the single-degenerate and double-degenerate scenarios. In the former, a white dwarf accretes hydrogen-rich material from a non-degenerate companion, and this material is then processed into carbon and oxygen via thermonuclear burning on the surface of the accretor \\citep{Whelan.Iben.1973}. In the DD scenario, a binary system goes through two separate white dwarf formation events (or -- in rare cases -- a single event in which both components become white dwarfs simultaneously, see \\citealt{Toonen.et.al.2012}), after which the binary loses angular momentum, until the two components merge to form a single white dwarf with a mass at or above that required for a type Ia supernova explosion \\citep{Iben.Tutukov.1984,Webbink.1984}. Both scenarios have been studied intensely, but due to the relative scarcity of type Ia supernova events in nearby galaxies direct observational evidence is lacking.  One promising way to distinguish between progenitor scenarios is by X-ray observations. The reason is that the thermonuclear burning of hydrogen-rich material on the surface of the white dwarf in the single-degenerate scenario is expected to be highly luminous in supersoft X-rays. This would mean that single-degenerate progenitors are so-called supersoft X-ray sources (SSSs), characterised by near-Eddington luminosities ($L_{\\mathrm{bol}} \\sim 4\\cdot 10^{37} - 10^{38}$ erg/s) and low effective temperatures ($kT_{\\mathrm{eff}}=$ 30-150 eV) \\citep{van.den.Heuvel.et.al.1992,Kahabka.van.den.Heuvel.1997}. This SSS state of a single-degenerate progenitor would presumably persist up to the point where the accreting white dwarf explodes as a type Ia supernovae. On the other hand, DD progenitors are not expected to be SSSs \\citep{Nielsen.et.al.2014a}. If pre-explosion observations of the positions of known type Ia supernovae can be obtained it may therefore be possible to establish if the progenitors are single-degenerate, accreting systems. As a consequence, a systematic search of the \\textit{Chandra} data archive is being undertaken, see \\citet{Voss.Nelemans.2008,Roelofs.et.al.2008,Nielsen.et.al.2012,Nielsen.et.al.2013b}, but so far, no unambiguous direct detections of supersoft X-ray emissions from type Ia supernova progenitors have been found. Moreover, the number of observed SSSs and the integrated soft X-ray luminosity of gas-poor galaxies both fall one to two orders of magnitude short of accounting for the expected number of massive, accreting white dwarfs, provided the single-degenerate scenario is the dominant contributor to the type Ia supernova rate \\citep{DiStefano.2010,Gilfanov.Bogdan.2010}. This constraint also holds if single-degenerate progenitors radiate at lower temperatures than that typical of most SSSs \\citep{Woods.Gilfanov.2013,Johansson.et.al.2014}.  On January 22 2014, the discovery of a new supernova (initially designated PSN\\_J09554214+6940260) in M82 was announced and categorised as a young type Ia \\citep{Cao.et.al.2014}. Shortly afterwards, it was re-designated SN2014J, and the exact position was given as 09:55:42.121,+69:40:25.88 \\citep{STSI.2014}. At a distance of $\\sim$3.5 Mpc \\citep{Dalcanton.et.al.2009}, this is the closest type Ia supernova since SN1972E more than 40 years ago. Since then, a number of sensitive satellite X-ray telescopes have been in operation, including \\textit{EXOSAT}, \\textit{Einstein}, \\textit{ROSAT}, \\textit{XMM-Newton} and \\textit{Chandra}. As a result, large amounts of archival X-ray data are now available. The relatively short distance to M82 raised the possibility that SN2014J might yield much-needed, strong observational constraints on the progenitor problem.  As a continuation of the systematic archival search mentioned above, the authors analysed the publicly available pre-explosion \\textit{Chandra} observations of the position of SN2014J to establish if an X-ray source was present. Additionally, we analysed the single available post-explosion observation of SN2014J to establish upper limits on the X-ray luminosity of the supernova itself. In Section \\ref{Sect:Data.Reduction} we presents our data analysis, and Section \\ref{Sect:Discussion} disusses the implications of our results. Section \\ref{Sect:Conclusion} concludes. This article represents an elaboration on the preliminary results reported by \\citet{Nielsen.et.al.2014b}, see also \\citet{Maksym.et.al.2014} ",
        "conclusions": "\\label{Sect:Conclusion} SN2014J in M82 is the nearest type Ia supernova in four decades, and the position of the progenitor system has been observed extensively by the \\textit{Chandra} X-ray Telescope prior to the supernova explosion.  We examined the totality of pre-explosion observations from the \\textit{Chandra} data archive and found no evidence of an X-ray progenitor for SN2014J. We calculated upper limits on the X-ray and bolometric luminosities of the progenitor (Tables \\ref{Table:Results}, \\ref{Table:Results_Xray_ULs}, Fig.\\ref{Fig:Teff_NH}). Our results firmly rule out the progenitor being a single-degenerate system with a photospheric radius comparable to the near-Chandrasekhar mass white dwarf radius and mass accretion rate in the range where stable nuclear burning can occur. This is in agreement with results obtained for several other nearby type Ia supernovae with pre-explosion images in the \\textit{Chandra} archive. Due to the presence of a large absorbing column we cannot rule out a low temperature supersoft X-ray progenitor with $kT_{\\mathrm{eff}} \\lesssim$ 80 eV.  We also examined the one existing \\textit{Chandra} observation made after the supernova explosion, and found no evidence of X-ray emission on the position of the supernova. We calculated upper limits on the X-ray luminosity of SN2014J (Table \\ref{Table:Results_Xray_ULs_post_SN}).  SN2014J is located in the center of a region of enhanced soft diffuse emission, roughly 200 pc across, which appears to be a supernova-inflated bubble around a region of recent star formation. The short cooling time of the gas, estimated from the X-ray spectral analysis, suggests a relatively young age of the bubble. If SN2014J is indeed located inside the bubble, it belongs to the prompt population of type Ia supernovae, with the delay time of the order of $\\sim 50$ Myrs."
    },
    "1402/1402.2291_arXiv.txt": {
        "abstract": "We use hydrodynamic simulations to investigate nonlinear gas responses to an imposed stellar spiral potential in disk galaxies. The gaseous medium is assumed to be infinitesimally thin, isothermal, and unmagnetized. We consider various spiral-arm models with differing strength and pattern speed. We find that the extent and shapes of gaseous arms as well as the related mass drift rate depend rather sensitively on the arm pattern speed. In models where the arm pattern is rotating slow, the gaseous arms extend across the corotation resonance (CR) all the way to the outer boundary, with a pitch angle slightly smaller than that of the stellar counterpart. In models with a fast rotating pattern, on the other hand, spiral shocks are much more tightly wound than the stellar arms, and cease to exist in the regions near and outside the CR where $\\Mperp/\\sin p_* \\simgt 25$--40, with $\\Mperp$ denoting the perpendicular Mach number of a rotating gas relative to the arms with pitch angle $p_*$. Inside the CR, the arms drive mass inflows at a rate of $\\sim 0.05$--$3.0\\Aunit$ to the central region, with larger values corresponding to stronger and slower arms. The contribution of the shock dissipation, external torque, and self-gravitational torque to the mass inflow is roughly 50\\%, 40\\%, and 10\\%, respectively. We demonstrate that the distributions of line-of-sight velocities and spiral-arm densities can be a useful diagnostic tool to distinguish if the spiral pattern is rotating fast or slow. ",
        "introduction": "Disk galaxies possess prominent non-axisymmetric features such as spiral arms and bars that have profound influences on galactic evolution in various ways (e.g., \\citealt{but96,kor04,but13,sel13} and references therein). Stellar spiral arms not only trigger or organize star formation in the outer parts of disk galaxies but also drive secular changes in the orbits of stars and gas clouds, redistributing the mass in the disks (e.g., \\citealt{lin64,lin66,too64,elm95,ber96,foy10}). Bar potentials also affect the mass redistributions in the inner parts and are responsible for the formation of various gaseous substructures such as dust lanes and nuclear rings (e.g., \\citealt{san76,ath92,hel94,pin95,but96,kim12a}). Understanding the gravitational interaction of the stellar potentials with a gaseous medium is therefore the first step to understand star formation, secular evolution, and morphological changes occurring in disk galaxies.  Among various secular processes, an angular momentum exchange between gas and a stellar pattern is particularly interesting since it leads to overall gas inflows or outflows in the radial direction. In barred galaxies, it has been well established that a non-axisymmetric torque exerted by a bar potential produces a pair of dust lanes in the gaseous medium, across which the gas loses angular momentum and falls radially inward to form a nuclear ring (e.g., \\citealt{ath92,pin95,eng97,pat00,mac04,ann05,tha09,kim12a,kim12}), potentially powering star burst activities in nuclear rings as well as fueling active galactic nuclei (e.g. \\citealt{shl90,reg99,kna00,lau04,jog05,hun08,van10}). In particular, \\citet{kim12b} demonstrated that the location of nuclear rings is determined not by the resonances but by the centrifugal barrier that the inflowing gas cannot overcome.  This suggests that nuclear rings are smaller in size in more strongly-barred galaxies, entirely consistent with the observational results of \\citet{com10}.  Compared to the cases with a bar potential, the effects of a spiral potential on the gaseous structures and radial mass inflow are relatively poorly understood. While interactions of stellar density waves with stars and the associated stellar heating and radial migration have been a subject of intense study (e.g., \\citealt{lyn72,don94,zha96,ath02,sel02,ros08,oh08,sel11,bru11}; see also \\citealt{kor04} and \\citealt{sel13} and references therein), only a few studies have explored angular momentum loss of gas due to the density waves (e.g., \\citealt{kal72,rob72,lub86,hop11}). For example, \\citet{rob72} considered non-self-gravitating galactic spiral shocks and showed that the damping timescale of stellar density waves due to the angular momentum exchange is of the order of $\\sim1\\Gyr$. \\citet{lub86} included the back reaction of the density waves on the gas density distribution, finding that the gas accretion rate due to the stellar pattern amounts to $\\Mdot\\sim -(0.2$--$0.4)\\Aunit$ for parameters representing the solar neighborhood in the Milky Way. These values of $\\Mdot$ are overall consistent with the mass inflows inferred from chemical modeling in the Milky Way \\citep{lac85} and also those in external galaxies based on gravitational torque analyses \\citep{haa09,gar09}. \\citet{hop11} used the epicycle approximation to derive an analytic expression for $\\Mdot$ due to a non-axisymmetric potential.  While the derivations of $\\Mdot$ by \\citet{lub86} and \\citet{hop11} are insightful, they utilized a few notable approximations. First, \\citet{lub86} considered local, tightly-wound waves in both the stellar and gaseous media and ignored the self-gravitational torque on the gas in evaluating $\\Mdot$. They also included shear viscosity to represent cloud collisions, which tends to smear out shock profiles and thus makes it difficult to isolate the sole effect of the shock (e.g., \\citealt{kim07,kim08}). On the other hand, \\citet{hop11} used the orbit crossing of test particles as a criterion for the shock formation, without considering the effects of gas pressure as well as the speed of incident flows relative to the pattern. Although \\citet{hop11} showed that their $\\Mdot$ is in good agreement with their numerical results for galaxies with a dominant bar-like potential, it is uncertain whether the same holds true for spiral galaxies in which the effects of thermal and ram pressures are more important than in barred galaxies.  Since dynamics involved with spiral arms is intrinsically nonlinear, it is desirable to run numerical hydrodynamic simulations in order to measure the mass drift rates driven by spiral arms properly. There have been numerous studies for gas responses to an imposed spiral potential, focusing on morphological changes of gaseous arms depending on the pattern speed (e.g., \\citealt{pat94,pat97,gom02,sly03,yan08,gom13}), formation of arm substructures such as branches, spurs, and feathers (e.g., \\citealt{cha03,wad04,she06,dob06}; see also local models of \\citealt{kim02,kim06}), or star formation occurring in spiral arms (e.g., \\citealt{she08,wad08,wad11,dob11}). While the arm-driven mass inflows might have affected the simulation outcomes in the work mentioned above, its rate has yet to be evaluated to assess its dynamical consequences on secular evolution quantitatively.  In this paper, we run global hydrodynamic simulations for gas evolution in galaxies with spiral potentials. We consider an infinitesimally-thin, uniform gaseous disk.  We take an isothermal equation of state for the gas, and ignore the effect of radiative cooling and heating, star formation, and magnetic fields. Two important parameters characterizing a spiral pattern are its angular frequency $\\Oms$ and strength $\\F$, which are difficult to constrain observationally. Thus we in this work vary $\\Oms$ and $\\F$ as free parameters to model spiral arms in various galactic situations, and study how $\\Mdot$ depends on them. We will also compare our numerical results with the analytic expression presented in \\citet{hop11}.  In addition to evaluating $\\Mdot$, our models are also useful to address important issues related to the spatial extent, structures, and pitch angles of gaseous arms in comparison with their stellar counterparts. While the theory for spiral density waves suggests that the \\emph{stellar} pattern extends up to the corotation resonance (CR) or to the outer Lindblad resonance (OLR) if it is in the linear-regime (e.g., \\citealt{too81,lin79,ber89a,ber89b,zha96}) and to the 4/1 resonance if it is strong enough to be nonlinear \\citep{con86,con88,pat91}, it is uncertain whether the termination of \\emph{gaseous} arms corresponds to the resonance radii. Moreover, \\citet{git04} showed using both semi-analytic and numerical approaches that gaseous arms are in general more tightly wound than the stellar arms. Although large uncertainties surround observational determinations of arm pitch angles, recent studies show that they are, statistically, slightly larger in the $I$- or $H$-band than in the $B$-band \\citep{sei06,dav12,mart12}, suggesting that the gaseous arms are likely more tightly wound than the stellar counterpart. We will show that the extent and pitch angles of gaseous arms are dependent somewhat sensitively upon the arm pattern speed. We will also show that the distributions of line-of-sight velocities in the projected galactic disk and density profiles of gaseous arms can be used to tell whether the observed arms are inside their CR or not.  The remainder of this paper is organized as follows. In Section \\ref{sec:model}, we describe the galaxy model and our choices of the model parameters, as well as the numerical method we use. In Section \\ref{sec:arm}, we present the simulation results on morphologies of spiral shocks. In Section \\ref{sec:mdot}, we measure the arm-induced mass drift rates as functions of the pattern speed and strength of the stellar arms. In Section \\ref{sec:los}, we present the distributions of line-of-sight velocities in the plane of sky and discuss how they can be used to obtain information on the arm pattern speed. In Section \\ref{sec:sum_dis}, we conclude with a summary and discussion of our results and their astronomical implications. ",
        "conclusions": "\\label{sec:sum_dis}  \\subsection{Summary}\\label{sec:sum}  We have presented the results of grid-based hydrodynamic simulations on spiral structures and radial mass drift in disk galaxies driven by spiral arms. The gaseous disk is assumed to be infinitesimally thin, unmagnetized, isothermal with the sound speed of $\\cs=10\\kms$, and initially uniform with surface density $\\Sigma_0=10\\Surf$. For the spiral arms, we impose a rigidly-rotating logarithmic gravitational potential with pitch angle $p_*=20\\degr$, strength $\\F$, and pattern speed $\\Oms$. To study the dependence of the shock structure and the mass drift rates on $\\F$ and $\\Oms$, we consider three types of models in which the arm is rotating fast at $\\Oms=30\\freq$, intermediately at $\\Oms=20\\freq$, or slow at $\\Oms=10\\freq$, which are referred to as the F, I, and S model, respectively. We also vary $\\F$ between 5 and 20\\%. The main results of this paper are summarize as follows.  1. \\emph{Extent of Spiral shocks.} -- The radial extent of spiral shocks depends rather sensitively on the arm pattern speed. In the F and I models, spiral shocks exist only up to $\\Rt\\sim 17\\kpc$ and $\\sim25\\kpc$, respectively, while the outer region is almost featureless other than weak trailing waves. This is because when equation (\\ref{eq:t_ratio}) is not fulfilled, gas perturbed by one arm does not have sufficient time to adjust itself to the imposed spiral potential before encountering the next arm. That is, the rapid rotation of the potential makes itself effectively smoothed considerably along the azimuthal direction, and gas moves almost ballistically \\citep{bak74}. In the S models, on the other hand, the whole disk satisfies $\\Mperp/\\sin p_*\\simlt 20$ and forms spiral shocks across the entire simulation domain. In these models with a slow pattern, spiral shocks are not terminated at the CR due to epicycle motions of perturbed gas elements. Since a gas element on its epicycle orbit achieves larger (smaller) amplitudes as it moves radially outward (inward), the spiral shocks produced inside the CR are able to extend all the way to the outer radial boundary, while those originally formed outside the CR extend only slightly inward of the CR.  As a consequence, the dense arm gas outside the CR is bounded by two spiral shocks in the S models.  2. \\emph{Relation between the pitch angle and shock strength.} -- In a quasi-steady state, stronger spiral shocks tend to form at farther downstream relative to the minima of the imposed spiral potential. Since $\\Mperp$ varies systematically with $R$, this makes the pitch angle $p_{\\rm gas}$ of the gaseous arms smaller than that of the stellar arms. In our models, the offset between $p_{\\rm gas}$ and $p_*$ amounts to $\\sim2\\degr$--$12\\degr$, and is larger for smaller $\\F$ since a deeper potential tends to have shocks closer to its minima. It is also larger for models with larger $\\Oms$ due to larger radial variations of $\\Mperp$.  Equation (\\ref{eq:delp}) gives our fits to $\\Delta p=p_* - p_{\\rm gas}$ against the peak shock density $\\Sigp$ of the gaseous spiral arms.  3. \\emph{Mass Drift} -- The non-axisymmetric spiral potential is an efficient means of angular momentum transport, initiating radial drift of gas that would otherwise be in circular motions. In our models, the radial mass drift is caused by a combination of three processes: angular momentum loss at spiral shocks, external gravitational torque, and self-gravitational torque. While self-gravitational torque is always negative, it is usually smaller than the other torques inside the CR.  On the other hand, the direction of the mass drift by the shock loss and external torque depends on the sign of $\\Omega-\\Oms$, such that it is radially inward inside the CR and outward outside the CR. The resulting mass inflow rate, averaged over $\\RILR\\simlt R\\simlt \\RCR$, is in the range $\\aMtot_{\\rm in}\\sim -(0.05$--$3.0)\\Aunit$, with a larger value corresponding to stronger and/or slower arms, as described by equation (\\ref{eq:Mtot2}). The shock loss and external spiral potential account for about 50\\% and 40\\% of the total, respectively.  4. \\emph{Line-of-Sight Velocity} -- Since the spiral arms cause streaming motions in the gas flows whose amplitudes depend on the arm pattern speed, the related line-of-sight velocity $\\Vlos$ and the density distribution across the arms can potentially provide information on the arm pattern speed. Gaseous arms located outside the CR have a larger density at larger $R$ along a radial cut. In this case, the gas in the arms rotates faster than the pattern due to the streaming motions, tending to make the locus of $\\Vlos=0$ bend in the same way as the direction of galaxy rotation. In contrast, gaseous arms located inside the CR have a larger density at smaller $R$, and the $\\Vlos=0$ curve near the arms bends in the opposite sense to the galaxy rotation (see Fig.\\ \\ref{fig-vlosmod}).  \\begin{figure*} \\centering \\includegraphics[angle=0,width=0.98\\textwidth]{fig_14.eps} \\caption{Logarithm of gas surface density at $t=1.0\\Gyr$ for self-gravitating models with pitch angle (a) $p_*=20\\degr$, (b) $p_*=30\\degr$, and (c) $p_*=44\\degr$.  The other parameters are the same as in Model I10G. In each panel, the dotted and solid circles draw the CR and OLR.  A ring-like structure forms near the position where the outer end of a gaseous arm curls radially in to be connected to the other arm. Colorbar labels ${\\rm log}(\\Sigma/\\Sigma_0)$. \\label{fig-res}} \\end{figure*}  \\subsection{Discussion}\\label{sec:dis}  Our result that the radial extent of gaseous arms depends on the arm pattern speed is overall consistent with the finding of \\citet{pat94} who showed using SPH simulations that gaseous spirals exist up to the CR when $\\Oms$ is large ($\\sim37.4\\freq$), while they extend to the end of the stellar spiral arms when $\\Oms$ is small ($\\sim12.5\\freq$). \\citet{pat94} further found that an oval ring forms, possibly by the resonance, near the OLR area.  A similar ring-like structure forms in the outer regions of Model I10G  at late time when the outer ends of trailing gaseous arms curve back radially inward to touch the other arms (see Fig.\\ \\ref{fig-sigt}). The main difference in the model parameters between our and their models is the pitch angle of the stellar potential. To explore the effect of $p_*$ on the position of ring-like structures that form, we have run two additional models with $p_*=30\\degr$ and $44\\degr$, while the other parameters remain identical to Model I10G.\\footnote{Models in \\citet{pat94} took $p_*=44\\degr$.} Figure \\ref{fig-res} plots the distributions of surface density at $t=1\\Gyr$ from these additional runs as well as Model I10G. In each panel, the dotted and solid circles mark the CR and OLR, respectively. An oval-like structure forms in all models, although its size tends to decrease with increasing $p_*$. Model I10G with $p_*=20\\degr$ has an oval-like structure well outside the OLR, while it is coincidentally at the OLR. Since the arm termination (and thus the ring formation) occurs at the radius where equation (\\ref{eq:t_ratio}) is satisfied, one can expect a smaller ring when $p_*$ is larger, fully consistent with our numerical results. This demonstrates that the formation of a ring-like structure near or beyond the OLR is not due to the resonance.  It is interesting to apply our results to the spatial extent of observed arms in the barred-spiral galaxy M83. In this galaxy, the spiral pattern rotates relatively rapidly at $\\Oms\\approx 45\\freq$ and has a pitch angle of $p_*\\approx 16\\degr$ (e.g., \\citealt{lor91,zim04}). The rotational velocity at the flat part is $\\sim180\\kms$ \\citep{lun04}, so that the CR is located at $\\sim4\\kpc$, corresponding to $\\sim3^\\prime$ at the distance of $4.5\\Mpc$ \\citep{thm03}. The radio data of \\citet{lun04} show that the gaseous arms are weaker at the CR than at the neighboring arms, similarly to our results shown in Figure \\ref{fig-mach}. The arms extend up to $\\sim 6^\\prime$, while there is a plenty of gas with $\\Sigma\\sim2\\Surf$ in the outer regions where the gaseous arms are absent \\citep{cro02,lun04}. Although the termination radius of the gaseous arms in M83 is close to the OLR, it is uncertain whether the OLR plays a central role in limiting the arm extent.  We note that the radius of $6^\\prime$ corresponds to $\\Mperp/\\sin p_*\\sim 23$--30 for the observed CO velocity dispersions of $\\cs\\sim7.8\\pm0.9\\kms$ \\citep{lun04}, suggesting that the idea of arm termination by too large $\\Mperp$ is not inconsistent with the observed gaseous arms in M83 with $\\F\\sim5$--10\\%.  To measure the mass drift rate unaffected by a radial density gradient for given $\\F$, we have employed simple disk models with radially constant $\\Sigma_0$ and $\\F$.  In models with a slow patten speed, this inevitably results in readily discernable spiral shocks all the way to the outer radial boundary. In reality, however, gas surface density in spiral galaxies appears to drop off exponentially or more rapidly (e.g., \\citealt{big12}). In addition, the stellar spiral potential is likely to become shallower with increasing $R$ beyond the CR (e.g., \\citealt{con86,con88,pat91}). Although arm-to-interam density contrasts are likely unchanged by the background density (especially when self-gravity is unimportant), small values of $\\Sigma_0$ and $\\F$ would make it difficult to detect gaseous spiral arms at large radii in real spiral galaxies.   The tendency of spiral shocks moving toward the upstream direction with increasing $R$ was reported by \\citet{git04}, and our results further show that the displacement of spiral shocks is larger when $\\Oms$ is larger. This is consistent with the results of \\citet{pat94} who found that gaseous arms are much tighter than the stellar pattern in their high-$\\Oms$ models. \\citet{git04} also noted that in addition to stellar and gaseous arms, there are star-forming arms traced by \\HII regions, all of which may have different pitch angles such that $p_* > p_{\\rm gas} > p_{\\rm SF}$ if the time offset between the gaseous and star-forming arms is independent of $R$. Indeed, \\citet{gro98} showed that the stellar arms traced by $K^\\prime$-band observations are more loosely wound than the optical arms for a sample of five galaxies. Although \\citet{dav12} more recently found that the arm pitch angles in the optical band are almost identical to those in the near-IR band within observational uncertainties (see also \\citealt{sei06}), their fitting results plotted in their Figure 13 still show that the latter is larger by $\\sim2\\degr$ than the former for $p_* \\sim 10\\degr$--$30\\degr$, consistent with our numerical results. \\citet{mart12} also reported that arm pitch angles are larger in the $H$-band than in the $B$-band. More accurate observational estimates are required to explore the dependence on $\\Oms$ of the offsets between arm pitch angles in different wavelength bands.  In application to the gas accretion in the Milky Way, \\citet{lub86} used a local model of spiral shocks by considering the reaction of stellar waves back to the gaseous self-gravity and calculated the mass inflows rate due to the spiral shocks and external gravitational torque.  They found that the total inflow rate is in the range $\\Mdot\\sim-(0.2$--$0.4)\\Aunit$ for $\\F=3\\%$, consistent with the extrapolation of our numerical results. They also found that the external potential is about three times more effective than the viscous torque, which is different from out finding that the effect of the spiral shocks in removing angular momentum is slightly larger than that of the external potential. This difference is presumably caused by the fact that \\citet{lub86} included physical viscosity accounting for cloud collisions. This smears out the shocks and moves the peak density toward downstream (see also \\citealt{kim08,kim10}), which tends to enhance the external torque and reduce the angular-momentum loss at the shock fronts.  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.48\\textwidth]{fig_15.eps} \\caption{Comparison of the mass inflow rate $\\Mshock$ (solid lines) by angular-momentum loss at spiral shocks in our models with the analytic estimate $\\Mdot_{\\rm HQ}$ of \\citet{hop11} (dashed lines) given by equation (\\ref{eq:HQ}) for the (a) S, (b), I, and (c) F models. Note that $\\Mdot_{\\rm HQ}$ agrees approximately with $\\Mshock$ only in the S models with small $\\F$, while it differs considerably from $\\Mshock$ for the small-$\\Oms$ with large $\\F$ and all the large-$\\Oms$ models.\\label{fig-hq}} \\end{figure}  \\citet{hop11} put forward a simple analytic model for angular momentum transport and related gas inflows driven by a non-axisymmetric stellar potential. By neglecting the effects of thermal pressure as well as the flow velocity relative to the potential, they considered orbit crossing of collisionless particles as a criterion for shock formation and derived the mass inflow rate \\begin{equation}\\label{eq:HQ} \\Mdot_{\\rm HQ} = -\\bar\\Sigma R^2 \\Omega \\F\\tan p_* {\\rm sign}(\\Omega-\\Oms) |f(\\zeta)|, \\end{equation} where $\\bar\\Sigma$ is the azimuthally-averaged gas density and $f(\\zeta)$ denotes the correction factor of order unity for the degree of the orbit crossing, $\\zeta$. They confirmed that equation (\\ref{eq:HQ}) is in good agreement with their numerical results when bar-like stellar modes dominate the perturbations. To check whether equation (\\ref{eq:HQ}) is also good for the cases dominated by a spiral potential, Figure \\ref{fig-hq} compares $\\Mdot_{\\rm HQ}$ (dashed lines) with $\\Mshock$ (solid lines) from our numerical simulations. Note that equation (\\ref{eq:HQ}) can be a good approximation for the mass inflows due to the shock loss \\emph{alone}, only for weak spiral shocks in the S models in which shock fronts are located very close to the potential minima. On the other hand, equation (\\ref{eq:HQ}) fails to describe the mass drift accurately for the slow-arm models with large $\\F$ and for the F models where spiral shocks are displaced significantly from the potential minima. In this case, the shock formation requires consideration of thermal pressure as well as incident velocities, which were neglected in \\citet{hop11}.  By analyzing Sloan Digital Sky Survey DR7 data, \\citet{com09} found that about 20\\% of 266 galaxies with measured bar strength host dust lanes. On the other hand, numerical simulations with only a bar potential without spiral arms show that dust lanes remain strong only for $0.1\\Gyr$ around the time when the bar potential achieves its full strength (e.g., \\citealt{kim12b,kim12}). This implies that barred galaxies with strong dust lanes should be either dynamically young or supplied with fresh gas. Our numerical results in this paper suggest that spiral arms can be efficient to transport gas from outside to the central region provided that the spiral arms have quite a low pattern speed so as to have a large CR radius.  In addition to enhancing the mass in the central region, mass inflows by spiral arms also may help to increase the rate of star formation occurring in the nuclear ring induced by a bar potential. \\citet{seo13} numerically found that without spiral arms or gas infall from halo, the star formation rate in nuclear rings decays to small values below $\\sim 1 \\Aunit$ after showing a strong burst that lasts only about $0.1\\Gyr$. This is in contrast to the claim that star formation in nuclear rings is a long-lived phenomenon ($\\sim1$--$2\\Gyr$), with multiple episodes (e.g., \\citealt{all06,sar07,van13}). In Model S10G, gas flows radially inward at a rate $\\Mtot \\sim -1.0\\Aunit$ at $R=5\\kpc$, which would fuel star formation on the galactic center when enough mass is accumulated to undergo gravitational collapse. It would be interesting to study how much star formation is enhanced in nuclear rings by an addition of outer spiral arms, which will direct our future research.  By employing a simple isothermal equation of state, our models produce spiral shocks with quite a large density contrast between arm and interarm regions.  For instance, \\citet{she07} reported that spiral arms seen in CO observations of M51 typically have an arm-to-interarm density contrast of $\\sim5$, while it is $\\sim 20$ in our models (see, e.g., Fig.\\ \\ref{fig-Phi}). Even considering the beam smearing of the CO data, the density contrast in our models seems to be larger than observed values. In real galaxies, there are many physical processes including magnetic fields and star formation that can affect spiral-shock gas dynamics considerably.  Magnetic fields that are pervasive in the interstellar medium (e.g., \\citealt{bec96,fle11}) can make the isothermal gas behave as if it were adiabatic with index of 2 for one-dimensional compression (e.g., \\citealt{shu92}), reducing the arm-peak density substantially (e.g., \\citealt{kim02,kos02,lee12}). Star formation and ensuing feedback occurring inside spiral arms is able to provide the arm gas with turbulent kinetic energy which not only tends to disperse dense gas but also triggers new star formation (e.g., \\citealt{mk04,mck07,she12,kim13}).  A reduced arm density may decrease the mass drift rate, while magnetic fields themselves can be a source of additional angular momentum transport via tension forces. Therefore, more realistic quantitative assessment of the mass drift rate caused by spiral arms requires consideration of these processes as well as radiative heating and cooling.  Finally, we discuss the distribution of gas density and $\\Vlos$ in the arms found in our models in comparison with observations. \\citet{reg02} presented a velocity map of CO emission from the barred-spiral galaxy NGC 3672, derived from Gaussian fits to the line profiles. Their Figure 6 shows that the observed distributions of $\\Vlos$ in the outer regions, especially the loci of $\\Vlos=0$, match those in Model S10G better than in Model F10G presented in Figure \\ref{fig-vlosmod}, suggesting that the spiral arms in NGC 3672 are located inside the their CR. A close inspection of Figure 4a of \\citet{reg02} (see also Figure 1c of \\citealt{che03}) reveals that the density profiles of the gaseous arms along a radial cut have a steeper gradient at the inner edge than the outer edge, which also hints that the spiral arms rotate slow. Since the CR of a strong bar is placed just outside the bar end, these all indicate that the spiral arms in NGC 3672 have a lower pattern speed than the bar. Indeed, \\citet{ran04} used the method of \\citet{tre84} to show that the bar in NGC 3627 has a pattern speed of $\\Omb=50^{+3}_{-8}\\freq$, while the southern extension of the western spiral arm has $\\Oms=23\\pm4\\freq$. The corresponding CR radii of the bar and arms are at $R=1.3^\\prime$ and $3.3^\\prime$, respectively, indicating that the observed spiral arms in \\citet{reg02} are located in between the CRs of the bar and spiral arms, consistent with kinematic features in our models."
    },
    "1402/1402.2258_arXiv.txt": {
        "abstract": "We consider the pertubative consistency of inflationary models with features with effective field theory methods. By estimating the size of one-loop contributions to the three-point function, we find the energy scale where their contribution is of the same order of the tree-level amplitude. It is well-known that beyond that scale, perturbative unitarity is lost and the theory is no more under theoretical control. Requiring that all the relevant energy scales of the problem are below this cutoff, we derive a strong upper bound on the sharpness of the feature, or equivalently on its characteristic time scale, which is independent on the amplitude of the feature itself. We point out that the sharp features which seem to provide better fits to the CMB power spectrum are already outside this bound, questioning the consistency of the models that predict them. ",
        "introduction": "The recent analysis of \\emph{Planck} data does not show any significant deviations from the simplest single-field slow-roll models of inflation, but some issues remain open. In particular, it seems that there could be a relative better fit to the curvature power spectrum if the possibility of small and rapid oscillations is taken into account \\cite{Planck_ps, Hu2013, Benetti2013, Achucarro2013}. Even though, from present data, it seems that there is not enough improvement in the fits to assess their statistical significance \\cite{Spergel2013a, Spergel2013b}, it is very interesting to study the models that could provide a primordial origin to these signatures. In this paper we will focus on the so-called ``models with features'', which are well-known in the literature \\cite{Starobinsky1992, Wang1999, Adams2001, Peiris2003, Covi2006, Chen2006b, Covi2007, Chen2008, Piazza2008, Dvorkin2009, Hazra2010, Adshead2011, Adshead2011b, Chen2011a, Arroja2011b, Arroja2012, Miranda2012}. Tipically, oscillations in the power spectrum are due to some features in the potential of the inflaton or in the speed of sound, which induce a temporary deviation from the slow-roll dynamics that comes back to the attractor solution in less than few efolds without ruining inflation. Very interestingly, features also induce non-Gaussianity, giving us the possiblity to constraint these models also with an other observable, the bispectrum \\cite{planck_bs}.  In this context, from a theoretical point of view, a useful tool can be the Effective Field Theory of Inflation (EFTI) \\cite{Creminelli2007, Weinberg2008}. Independently of the mechanism that could produce the features, their effects on observables can equivalently be described through the time-dependent coefficients of the effective action of the Goldstone boson that non-linearly realizes time-diffeomorphisms, \\begin{eqnarray}\\label{S_eft} S = & \\displaystyle \\int \\dif^4 x \\sqrt{-g} & \\left[-M_{Pl}^2\\dot{H}(t+\\pi)\\left(\\dot{\\pi}^2-\\frac{{(\\partial_i \\pi)}^2}{a^2}\\right) \\right.+\\nonumber \\\\ & & +\\,2M_2^4(t+\\pi)\\left(\\dot{\\pi}^2 + \\dot{\\pi}^3 -\\dot{\\pi} \\frac{{(\\partial_i \\pi)}^2}{a^2}\\right)- \\nonumber \\\\ & & \\left. \\frac{4}{3}M_3^4(t+\\pi)\\dot{\\pi}^3+  .\\,.\\,. \\right]\\; , \\end{eqnarray} where $\\pi$ is related to the curvature perturbation $\\zeta$ by $\\zeta=-H\\pi$. Consider, for example, a time dependence of the form \\begin{equation}\\label{H_feat} \\dot{H}(t)=\\dot{H}_0(t)\\left[1+\\epsilon_{step}\\,F\\left(\\frac{t-t_f}{b}\\right)\\right] \\; , \\end{equation} where the function $F$ represent a step centered in $t_f$ with a height $\\epsilon_{step}$ and a characteristic width $\\Delta t = b$. It has been shown that, when inserted into the effective action \\eqref{S_eft}, one obtains the predicted damped oscillations in the power spectrum for models with features in the potential, simply neglecting the $M_n$ coefficients \\cite{Cannone2013}. Moreover the EFTI approach allow us to go beyond the standard scenario and generalize features in a model independent way to the speed of sound, $c_s$, and any other coefficients. Besides the spectrum, from the Taylor expansion \\begin{equation} \\dot{H}(t+\\pi)=\\dot{H}(t)+\\ddot{H}(t)\\pi+ ... \\; , \\end{equation} we get new interactions as, for example, the cubic term: \\begin{equation} \\mathcal{L}_3\\ni\\,-M_{Pl}^2\\ddot{H}(t)\\,\\pi\\dot{\\pi}^2 \\; , \\end{equation} which gives the largest contribution to the three-point function in the case of sharp feature. The resulting bispectrum is not scale-invariant and will be peaked for those modes that at the time of the feature, $t_f$, have an energy comparable to the inverse of the characteristic time scale, $b$, of the feature \\cite{Cannone2013}. If we define a sharpness parameter \\begin{equation} \\beta=\\frac{1}{bH} \\; , \\end{equation} as the ratio between the energy $1/b$ and the Hubble scale $H$, it can be shown that the modes that are most affected are more inside horizon as the feature becomes sharper and sharper \\cite{Chen2008, Chen2011a, Adshead2011}. Moreover, as the amplitude at the peak grows quadratically with the sharpness, this can become the major source of non-Gaussianity and be possibly seen in future data analyses. However, being $\\beta$ unconstrained, non-linearities could also be too large, not only for observations, but also for the theretical consistency of the models. ",
        "conclusions": "In this note, we considered the perturbative consistency of inflationary models with features by means of effective field theory methods. By estimating the size of loop contributions to the $n$-point functions and comparing them with the tree-level computation, one can identify the maximal energy scale at which the theory is unitary and perturbativity is safe. Then, by requiring that all the relevant energy scale of the physics we are interested in are below this UV cutoff, we can derive bounds on the parameters of the models. While in the standard slow-roll models of inflation, the only relevant energy scale is $H$, when features (or resonances) are present there is a new energy scale $E\\simeq1/\\Delta t$ corresponding to the inverse of the characteristic time-scale of the interaction. In the case of feature models, that we have studied here, we estimated the size of one-loop contributions to the three-point functions and compared them to the tree-level expectation. Our main result is that \\emph{there is a very strong upper bound on the sharpness $\\beta=1/\\Delta tH$ of the feature, eq. \\eqref{beta_bound_cs}, beyond which the unitarity of the theory is lost}. The result is general and valid for any feature either in the slow-roll parameters, the speed of sound or any other coefficient in the effective action for scalar perturbations during inflation. We point out that this bound is relevant even on observational grounds, as current best-fits of the CMB power spectrum seem to prefer sharp feature models with a sharpness $\\beta\\simeq300$, which is already outside our estimate of the unitarity bound."
    },
    "1402/1402.5364_arXiv.txt": {
        "abstract": "Recent observations by the Fermi satellite suggest that a photosphere emission component is contributing to the observed spectrum of many GRBs. One important question is whether the photosphere component can interpret the typical ``Band'' function of GRBs with a typical low energy photon spectral index  $\\alpha \\sim -1$. We perform a detailed study of the photosphere emission spectrum by progressively introducing several physical ingredients previously not fully incorporated, including the probability distribution of the location of a dynamically evolving photosphere, superposition of emission from an equal-arrival-time ``volume'' in a continuous wind, the evolution of optical depth of a wind with finite but evolving outer boundary, as well as the effect of different top-hat wind luminosity ($L_w$) profiles. By assuming a co-moving blackbody spectrum emerging from the photosphere, we find that for an outflow with a constant or increasing $L_w$, the low-energy spectrum below the peak energy ($E_{p}$), can be modified to $F_\\nu \\sim \\nu^{1.5}$ ($\\alpha \\sim +0.5$). A softer ($-1<\\alpha<+0.5$) or flat ($\\alpha=-1$) spectrum can be obtained during the $L_w$ decreasing phase or high-latitude-emission-dominated phase. We also study the evolution of $E_{p}$ as a function of wind and photosphere luminosity in this photosphere model. An $E_p-L$ tracking pattern can be reproduced if a certain positive dependence between the dimensionless entropy $\\eta$ and $L_w$ is introduced. However, the hard-to-soft evolution pattern cannot be reproduced unless a contrived condition is invoked. In order to interpret the Band spectrum, a more complicated photosphere model or a different energy dissipation and radiation mechanism are needed. ",
        "introduction": "}  Despite more than 40 years of observations of gamma-ray bursts (GRBs), the radiation mechanism during the prompt emission phase is still being debated. Observationally the time-resolved spectra are usually characterized by a smoothly-joint broken power law, known as the ``Band'' function \\citep{band93}. The typical value of the low-energy photon spectral index $\\alpha$ (the one below the peak energy $E_p$) is around -1 \\citep{preece00,zhang11,nava11}\\footnote{Some GRBs show a harder $\\alpha$ value in the time resolved spectra at the beginning or around the peak of the light curve \\citep{ghirlanda02,ghirlanda03}. In some other cases, a blackbody component can play a dominant role \\citep{ryde10,zhang11,ghirlanda13}.}. This value is not straightwardly predicted in available models. In general there are two broad categories of GRB prompt emission models, one invoking a Comptonized quasi-thermal emission from the photosphere of the outflow \\citep{thompson94,ghisellini99,meszarosrees00,meszaros02,rees05,peer06,thompson07,peer08,giannios08,beloborodov10,lazzati10,ioka10,toma11,peer11,mizuta11,lundman13,lazzati13,ruffini13}, the other invoking a non-thermal (synchrotron or synchrotron-self-Compton) mechanism in the optically thin region \\citep{meszaros94,tavani96,daigne98,lloyd00,zhangmeszaros02c,kumarpanaitescu08,zhangpeer09,wang09,kumarnarayan09,daigne11,zhangyan11,zhang12,veres12b,burgess11,uhm13}. Recent observations with Fermi \\citep{ryde10,zhang11,guiriec11,axelsson12,guiriec13} confirmed the pre-Fermi suggestion \\citep{ryde05,ryde09} that the observed GRB spectra sometimes include superposition of a quasi-thermal component with a non-thermal component\\footnote{Observational constraints demand that non-thermal component cannot be a single power law extending to much lower energies \\citep{ghirlanda07}.}, suggesting that the photosphere emission indeed contributes to the observed GRB spectra of at least some GRBs. The question is how much the photosphere emission contributes to the observed spectra. There are two different opinions. According to the first opinion, only the quasi-thermal component identified in some GRBs is of the photosphere origin \\citep[e.g.][]{ryde10,zhang11,guiriec11,axelsson12,guiriec13}, while the main ``Band'' component is non-thermal emission from an optically thin region. The second, more optimistic opinion suggests that the main Band component is of a photosphere origin, and the observed $E_p$ is defined by the photosphere temperature \\citep[e.g.][]{rees05,thompson07,beloborodov10,lazzati10,peer11,lundman13,lazzati13}.  Within this second scenario, one needs to broaden the quasi-thermal spectrum to a mimic the observed ``Band'' function. The high-energy photon index can be easily interpreted by introducing dissipation of energy near photosphere so that non-thermal electrons are accelerated to upscatter the seed quasi-thermal photons \\citep[e.g.][]{lazzati10}. However, the main difficulty is to account for the observed low-energy photon index $\\alpha \\sim -1$, since the predicted value is much harder than this value \\citep[e.g.][]{beloborodov10}. It has been speculated that geometric smearing and temporal smearing may soften the spectrum to make $\\alpha$ closer to -1, but no thorough study has been carried out.  Another interesting observational feature is the $E_p$ evolution in GRBs. Observtionally two patterns are identified: hard-to-soft evolution or $E_p$-intensity tracking \\citep{liang96,ford95,ghirlanda10,ghirlanda11a,ghirlanda11b,lu10,lu12} across a broad GRB pulse \\citep[or the ``slow'' variability component,][]{gao12}. Various data analyses suggest that both patterns co-exist, sometimes in a same burst. For short GRBs, the tracking behavior is most common, while for long GRBs, especially for the first broad pulse, the hard-to-soft evolution is relatively more common \\citep[e.g.][]{preece14}. Simulations show that in long GRBs, tracking pulses after the first pulse could be due to a superposition of multiple hard-to-soft pulses \\citep{lu12}, and it has been argued that the hard-to-soft evolution pattern may be ubiquitous among long GRBs \\citep{hakkila11,hakkila14}. It is of great interest to see what $E_p$ evolution pattern the photosphere model predicts.  In this paper, we investigate the predictions of $\\alpha$ values and $E_p$ evolutions within the simplest photosphere model (a co-moving blackbody spectrum and a uniform jet) by fully treating the geometrical and temporal smearing effects. In Sect.2, we describe our methodology, in particular, improvements upon previous work. We then present calculations of photosphere spectra and $E_p$ evolution patterns in progressively more complicated models in Sect.3. The conclusions are drawns in Sect.4. ",
        "conclusions": "In this paper, we have developed a sophisticated method to calculate quasi-thermal GRB photosphere spectra numerically by introducing several improvements on previous treatments \\citep{peer08,peer11}. The new ingredients introduced in this paper include: the probability distribution of the location of a dynamically evolving photosphere, superposition of emission from an equal-arrival-time ``volume'' in a continuous wind, the evolution of optical depth of a wind with finite but evolving outer boundary, as well as the effect of different wind luminosity profiles. By assuming a co-moving blackbody spectrum emerging from the photosphere and a top-hat jet profile, we address how these effects modify the observed spectra from blackbody. The following robust conclusions are drawn: 1. For an outflow with constant or increasing wind luminosity, the low-energy spectrum below $E_{p}$ can be modified to $F_\\nu \\sim \\nu^{1.5}$, corresponding to a low-energy photon index $\\alpha \\sim +0.5$. Introducing temporal smearing does not change $\\alpha$ significantly. 2. A softer spectrum can be obtained during the phase of decreasing wind luminosity with time, and a flat spectrum $F_\\nu \\sim \\nu^0$ ($\\alpha=-1$) can be obtained only when the spectrum is high-latitude emission dominated. However, since the photosphere radius is small, the flux drops very rapidly shortly after the wind terminates. 3. Depending on how $\\eta$ is related to $L_w$, this model can give both negative or positive $E_p - L_w$ correlation. The observed ``hard-to-soft'' evolution of $E_p$ across broad pulses of seconds duration \\citep{lu12} cannot be interpreted with this simple photosphere model, unless an unknown contrived physical condition to switch the index $m$ at the pulse peak is invoked. The intensity tracking patterns as observed in some broad pulses \\citep{lu12} can be accounted for this model if $\\eta \\propto L_w^m$ with $m>5/32$.  The results presented here suggest that the observed dominant spectral component, the so-called ``Band-function'' \\citep{band93} component, is not easy to interpret by this simplest photosphere model. The predicted low energy spectral index $\\alpha = +0.5$) is too hard compared with the typical observed value ($\\alpha = -1$), and the widely observed ``hard-to-soft'' $E_p$ evolution across broad pulses cannot be accounted for unless a contrived condition is invoked. In order to naturally interpret GRB spectra within the framework of the photosphere model, more complicated factors have to be considered. One possibility is to introduce energy dissipation (e.g. proton-neutron collisions, internal shocks, or magnetic reconnections) and particle heating around the photosphere region. Such a dissipative photosphere model can naturally account for a high energy tail through Compton scattering, but could not significantly modify the low-energy spectral index from $\\alpha \\sim +0.5$. \\cite{vurm11} introduced a synchrotron emission component, which peaks below the quasi-thermal component to make the ``effective'' low-energy spectral index softer. In order to make this synchrotron + quasi-thermal spectrum mimic a Band function as observed, the outflow magnetization parameter has to fall into a narrow range. Recently, \\cite{thompson13} invoked a magnetically dominated, low baryon-loading outflow, and modified the low-energy spectral index through the contribution from electron-positron pairs. Several authors pointed out the contrived conditions for the dissipative photosphere models to produce a single-component spectrum \\citep{vurm13,asano13,kumar13}. Another possibility to soften the spectrum below $E_p$ is to introduce a structured jet. \\cite{lundman13} showed that $\\alpha \\sim -1$ can be reproduced given that the GRB jets have a near constant $L_w$ but a structured Lorentz factor profile with angle. This can enhance the high-latitude contribution (large $1/\\Gamma$ cone at high-latitudes) to raise flux in large angles. For more general structured jets where both $L_w$ and $\\Gamma$ follow a certain angular profile \\citep[e.g.][]{meszaros98,zhangmeszaros02b,rossi02,zhang04}, the $\\alpha$ value would not be very different from what is calculated in this paper. In all these models, it is unclear how the ``hard-to-soft'' $E_p$ evolution commonly observed in many GRB pulses can be accounted for.  Alternatively, the main Band-component in the GRB spectra could arise from an optically-thin region well above the photosphere due to synchrotron radiation. \\cite{uhm13} recently showed that if the emission radius is large enough, the fast cooling problem for synchrotron radiation is alleviated, and $\\alpha \\sim -1$ can be reproduced in a moderately fast cooling regime. The hard-to-soft $E_p$ evolution pattern is a natural prediction in this model, since the outflow streams from small-radii where magnetic fields are stronger to large-radii where magnetic fields are weaker. Alternatively, the Band component may be interpreted as slow-cooling or slow-heating synchrotron emission in internal shocks where magnetic field strengths decays rapidly behind the shock \\citep{peer06b,asano09,zhao14}.  Recent Fermi observations revealed a quasi-thermal component superposed on the main Band component in a growing population of GRBs \\citep{ryde10,zhang11,guiriec11,axelsson12,guiriec13}. The spectral shape in our calculated photosphere emission is consistent with what is observed, suggesting that that component is very likely the photosphere emission from the GRB outflow \\citep{peer12}. This component is typically weaker than what is predicted in the standard fireball-internal-shock model, so that a certain degree of magnetization is needed for the outflow \\citep{zhangpeer09}. Within this picture, the non-thermal emission region in the optically-thin zone could be the internal shock region only if the magnetization parameter already falls below unity at the internal shock radius \\citep{daigne11}. It is possible that the outflow is still moderately magnetically dominated in the large zone. In this case, efficient GRB emission is possible due to internal-collision-induced magnetic reconnection and turbulence (ICMART) \\citep{zhangyan11,zz14}.    \\medskip We thank Xue-Feng Wu, Asaf Pe'er, Andrei M. Beloborodov, Z. Lucas Uhm, He Gao, Wei-Hua Lei, Hou-Jun L\\\"u, and Bin-Bin Zhang for helpful discussion or comments, and an anonymous referee for helpful suggestions. This work is partially supported by NASA under grant NNX10AD48G.    \\appendix     Plugging in $P(\\nu,T)$, $\\hat P(r,\\Omega)$ in Equation (\\ref{eq:f_nu_2_3_1}), one gets \\begin{equation} \\begin{array}{cl} \\hat F_{\\rm \\nu}(\\nu,t,\\hat t) & =  {\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int \\int \\hat P(r,\\Omega)\\cdot { n_\\gamma(\\nu,T) \\over 16\\pi({kT \\over hc})^3\\cdot \\zeta(3)}\\cdot h\\nu   \\cdot \\delta \\left(t - \\hat t - ({r u \\over \\beta c}-t_0) \\right)d\\Omega dr\\\\ \\\\ & ={\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int_{\\rm r_{min}}^{\\rm r_{max}} \\int \\int_0^{\\rm 2\\pi} {\\sigma_{\\rm T} n' \\Gamma {{\\cal D}^2 \\over 4\\pi} e^{-\\rm \\tau(r,\\mu,r_{\\rm out})} \\over A} \\cdot { n_\\gamma(\\nu,T) \\over 16\\pi({kT \\over hc})^3\\cdot \\zeta(3)}\\cdot h\\nu   \\cdot \\delta \\left(t - \\hat t - ({r u \\over \\beta c}-t_0) \\right)d(-\\mu) d\\phi dr\\\\ \\\\ & ={\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int_{\\rm r_{min}}^{\\rm r_{max}} \\int {\\sigma_{\\rm T} n' \\Gamma {{\\cal D}^2 \\over 2} e^{-\\rm \\tau(r,\\mu,r_{\\rm out})} \\over A} \\cdot { n_\\gamma(\\nu,T) \\over 16\\pi({kT \\over hc})^3\\cdot \\zeta(3)}\\cdot h\\nu   \\cdot \\delta \\left(t - \\hat t - ({r u \\over \\beta c}-t_0) \\right)d(-\\mu) dr\\\\ \\\\ & ={\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int_{\\rm r_{min}}^{\\rm r_{max}} \\int {\\sigma_{\\rm T} n' \\Gamma {1 \\over 2 \\beta (\\Gamma u)^2} e^{-\\rm \\tau(r,\\mu,r_{\\rm out})} \\over A} \\cdot { {8 \\pi {\\nu}^2 \\over c^3} {1 \\over {\\rm exp} ({h\\nu \\over kT} )-1} \\over 16\\pi({kT \\over hc})^3\\cdot \\zeta(3)}\\cdot h\\nu   \\cdot \\delta \\left(u - { \\beta c (t - \\hat t + t_0) \\over r} \\right){\\beta c \\over r}du dr ~~~~(u=1-\\beta \\mu)\\\\ \\\\ & ={\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int_{\\rm r_{min}}^{\\rm r_{max}} \\int {\\sigma_{\\rm T} n' \\Gamma {1 \\over 2 \\beta (\\Gamma u)^2} e^{-\\rm \\tau(r,\\mu,r_{\\rm out})} \\over A} \\cdot { {8 \\pi {\\nu}^2 \\over c^3} ( {\\rm exp} ({h\\nu \\Gamma u \\over kT^\\prime} )-1)^{-1} \\over 16\\pi({kT^\\prime \\over hc \\Gamma u})^3\\cdot \\zeta(3)}\\cdot h\\nu   \\cdot \\delta \\left(u - { \\beta c (t - \\hat t + t_0) \\over r} \\right){\\beta c \\over r}du dr ~~~~(T={T^\\prime \\over \\Gamma u})\\\\ \\\\ & ={\\dot N_0 \\over 4 \\pi d_{\\rm L}^2} \\int_{\\rm r_{min}}^{\\rm r_{max}} {\\sigma_{\\rm T} n' \\Gamma {r^2 \\over 2 \\beta (\\Gamma \\beta c (t - \\hat t + t_0))^2} e^{-\\rm \\tau(r,\\mu,r_{\\rm out})} \\over A} \\cdot { {2 {\\nu}^2 \\over c^3} ({ {\\rm exp} ({h\\nu \\Gamma  \\beta c (t - \\hat t+t_0) \\over kT^\\prime r} )-1})^{-1} \\over 4({kT^\\prime r \\over hc \\Gamma \\beta c (t - \\hat t+t_0)})^3\\cdot \\zeta(3)}\\cdot h\\nu {\\beta c \\over r} dr.\\\\ \\\\ \\end{array} \\label{eq:f_nu2} \\end{equation}   The limits of integration can be calculated from the formula of equal arrival time surface, $t - \\hat t = {r u \\over \\beta c}$, i.e.  $r={\\beta c (t - \\hat t) \\over u}$. With $\\theta_{\\rm min}=0$ and $\\theta_{\\rm max}=\\pi/2$, we get $r_{\\rm min}={\\rm max}[\\beta c (t - \\hat t),r_0]$ and $r_{\\rm max}={\\rm max}[{\\beta c (t - \\hat t) \\over 1- \\beta},r_0]={\\rm max}[\\Gamma^2 (1+\\beta) \\beta c (t - \\hat t),r_0]$."
    },
    "1402/1402.1771_arXiv.txt": {
        "abstract": "We describe the Carnegie-Spitzer-IMACS (CSI) Survey, a wide-field, near-IR selected spectrophotometric redshift survey with the Inamori Magellan Areal Camera and Spectrograph (IMACS) on Magellan-Baade. By defining a flux-limited sample of galaxies in Spitzer \\emph{IRAC} $3.6\\mu$m imaging of SWIRE fields, the CSI Survey efficiently traces the stellar mass of average galaxies to $z\\sim 1.5$. This first paper provides an overview of the survey selection, observations, processing of the photometry and spectrophotometry. We also describe the processing of the data: new methods of fitting synthetic templates of spectral energy distributions are used to derive redshifts, stellar masses, emission line luminosities, and coarse information on recent star-formation. Our unique methodology for analyzing low-dispersion spectra taken with multilayer prisms in \\emph{IMACS}, combined with panchromatic photometry from the ultraviolet to the IR, has yielded high quality redshifts for 43,347 galaxies in our first 5.3 degs$^2$ of the SWIRE XMM-LSS field. We use three different approaches to estimate our redshift errors and find robust agreement. Over the full range of $3.6\\mu$m fluxes of our selection, we find typical redshift uncertainties of $\\sigma_z/(1+z) \\la 0.015$. In comparisons with previously published spectroscopic redshifts we find scatters of $\\sigma_z/(1+z) = 0.011$ for galaxies at $0.7\\le z\\le 0.9$, and $\\sigma_z/(1+z) = 0.014$ for galaxies at $0.9\\le z\\le 1.2$. For galaxies brighter and fainter than $i=23$ mag, we find $\\sigma_z/(1+z) = 0.008$ and $\\sigma_z/(1+z) = 0.022$, respectively. Notably, our low-dispersion spectroscopy and analysis yields comparable redshift uncertainties and success rates for both red and blue galaxies, largely eliminating color-based systematics that can seriously bias observed dependencies of galaxy evolution on environment. ",
        "introduction": "\\label{sec:intro}  \\noindent Understanding the evolution  of galaxies and large scale structure remains a fundamental challenge in astrophysics.  Many ambitious galaxy surveys have been carried out to address this problem, but limited time on large telescopes results in a classic problem: sky coverage, depth, and spectral resolution -- choose two.  For example, the very-wide-area Sloan Digital Sky Survey (SDSS)  provides a wealth of spectral information for galaxies over a cosmologically significant volume, but its modest depth limits the SDSS to the relatively local, `modern' universe.  Conversely, the Hubble Ultra-Deep Field photometric survey probes deep into cosmic time --- to the epoch of reionization, but only over a tiny volume whose small population of galaxies leads to large uncertainties on their physical properties. While the union of such surveys has begun to paint a coherent picture of galaxy growth, significant patches of blank canvas limit our ability to fully describe how environmental processes, the infall of gas that fuels star formation, galaxy mergers and acquisitions, and feedback, all shape the evolution of galaxies over cosmic time.  The Carnegie-Spitzer-IMACS Survey (CSI) has been designed to address one of the most dramatic and least understood features of galaxy evolution --- the remarkably rapid decline in cosmic star formation since $z\\sim 1.5$.  It is during this extended epoch of galaxy maturation that galaxy groups and clusters have also emerged as a conspicuous feature of the landscape.  The CSI Survey is uniquely able to link together the evolution of individual galaxies with these features of large-scale structure growth.  In our ambitious spectrophotometric redshift survey of distant galaxies, we strike a balance between the aforementioned three factors: (1) high completeness to moderate redshift ($z\\simlt 1.5$), (2) spectral resolution intermediate between conventional photometric and spectroscopic surveys (combining the efficiency of imaging surveys with a spectral resolution high enough to resolve large-scale structure and prominent emission-lines); and (3)  an unprecedented area of 15 deg$^2$ for a $z\\simgt 1$ survey.  These give the CSI Survey a volume comparable to the SDSS, with a selection method that efficiently traces stellar mass over 2/3 the age of the universe ($0.4<z<1.5$) --- spanning the critical redshift range where cosmic star formation precipitously drops, and groups and clusters become prominent. As the redshift survey with the largest unbiased volume at $z=1$, CSI will allow us to comprehensively address the interplay between environment, galaxy mass buildup, and star formation at these redshifts.  The innovation of moving to low resolution for cosmologically interesting problems of galaxy evolution was instigated by the PRIMUS collaboration \\citep[e.g.][]{coil2011,cool2013}. They designed a low dispersion prism to perform an optically selected redshift survey of several legacy fields and have used their sample of 140,000 galaxies to constrain the evolution of galaxy stellar mass functions \\citep{moustakas2013}, and to study AGN activity \\citep{mendez2013} and it's relation to galaxy evolution \\citep{aird2012,aird2013}.  In this, the sample definition and basic analysis paper of the CSI Survey, we describe the design of the project, with particular focus on the flux limits of the selection, the data processing and the spectral energy distribution (SED) fitting with its attendant determination of redshifts and redshift errors. We apply our methodologies to the first batch of data in 5.3 degs$^2$ of the SWIRE XMM-LSS field, and provide an overview of some basic properties of galaxies over the past 9 Gyr, with an emphasis on the gains made by selecting galaxies by stellar mass instead of by their rest-frame UV light. ",
        "conclusions": "\\label{sec:general}  We conclude our discussion of the survey by presenting a broad overview of the CSI sample. Using the sample directly with the completeness functions derived earlier, we can estimate number densities of galaxies as a function of redshift and begin looking at galaxy quiescence and star formation activity. As we do so, the CSI selection limits are highlighted in order to give the reader a sense of the depth of the sample in luminosity, stellar mass, and redshift.  Figure \\ref{fig:nz} (top) shows the distribution of high quality CSI redshifts ($q\\in\\{2,3\\}$) in the CFHTLS XMM field. The decrease in $N(z)$ above $z=0.8$ is due in part to the declining success rates and completeness at fainter optical magnitudes (as shown earlier). Figure \\ref{fig:nz} (bottom) dissects the sample into two stellar mass ranges, $1<M/M*<2$ (orange), and $1/2<M/M_*<1$ (blue). Our completeness functions derived earlier allow us to straightforwardly correct the observed numbers of galaxies in CSI to accurate number densities down to magnitude limits of $[3.6]=21$ mag and $i=24.5$ mag. At stellar masses of $M=M*$ and $M=M*/2$, we estimate that the magnitude limits impose selection biases at $z\\simgt 1.45$ and $z\\simgt 1.2$, respectively; these limits are shown as hatched regions with the respective color. Above these redshifts, the $i$-band and \\emph{IRAC} flux limits explicitly exclude galaxies from the selection and no statistically meaningful statements about the galaxy census can be made. However, below these redshifts these subsamples remain statistically correctable.  Knowing that our samples are relatively free from bias down to masses of $M>M*/2$ back to $z=1.2$, we apply completeness corrections and convert our $N(z)$ figures to estimates of the comoving number densities and stellar mass densities of galaxies as functions of redshift. The resulting CSI number densities are plotted using the thick lines in Figure \\ref{fig:phiz}. The number and mass densities are shown for two ranges of galaxy stellar mass, orange for $1<M/M*<2$, and dark blue for $1/2<M/M_*<1$. Within these redshift bins the estimated cosmic variance of $\\sim 15\\%$ dominates the uncertainties ($\\sqrt{N}/N$ is small), and this uncertainty is shown by the hatched regions surrounding the thick orange and blue lines. The hatched regions shown at high redshift indicate redshifts where our faint magnitude limits impinge on these subsamples and the \\emph{IRAC} and $i$-band selections begin to exclude the oldest, passive galaxies. Down to the limits of CSI, Figure \\ref{fig:phiz} reinforces what previous deep surveys covering smaller fields-of-view have found:  a factor of two increase in the number densities of normal galaxies since $z=1$ \\citep[e.g.,][]{drory2009,marchesini2009,ilbert2010,moustakas2013}.   \\begin{figure*}[ht] \\centerline{ \\includegraphics[width=5.0in]{f22.pdf} } \\caption{Restframe $UVJ$ diagrams \\citep[e.g.][]{williams2009} for galaxies in CSI in several bins of redshift and stellar mass. The original \\cite{williams2009} $UVJ$ selection boundaries for quiescent galaxies are shown in black, superimposed on grayscale maps of the space density of galaxies at a given $U-V$ and $V-J$ color, mass, and redshift. The grayscales are fixed at a given mass scale to a range defined by our lowest redshift bin, to help illustrate the change in number density at high redshift at fixed mass. A green vector represents 1 mag of dust attenuation in the $V$-band for the \\cite{calzetti2000} extinction law. Blue contours highlight those galaxies with specific star formation rates $SSFR>2\\times 10^{-10}$ yr$^{-1}$ (at $z=0.1$, evolving to $5.5\\times 10^{-10}$ yr$^{-1}$ at $z=1$). These were derived using H$\\alpha$ or [OII]3727\\AA, depending on the redshift, corrected for dust extinction using $A_V$ from the burst components of our SED fitting, and applying the conversion from \\cite{kennicutt1998} or \\cite{kewley2004}. The red polygon shows the locus of galaxies with $SSFR < 10^{-11}$ yr$^{-1}$ (at $z=0.1$, evolving to $2.8\\times 10^{-10}$ yr$^{-1}$ at $z=1$) with a range of metallicities from 1/10th to twice solar, and $0\\le A_V\\le 0.5$ mag of dust attenuation. We use these polygons to define a selection for quiescent galaxies in the following figures. Note that for galaxies with stellar masses $M<M*/2$ at $z\\simgt 1.2$, the flux limits of CSI begin to select against quiescent galaxies, with these selection limits imposing on quiescent galaxies at lower masses at lower redshifts (see, e.g., Figure \\ref{fig:masslimits_lim}). \\label{fig:UVJ}} \\end{figure*}  \\begin{figure}[ht] \\centerline{ \\includegraphics[width=2.5in]{f23.pdf} } \\caption{The distribution of our $3.6\\mu$m-selected galaxies in restframe $B$-band magnitude as a function of redshift. The green and blue lines indicate the approximate loci for old stellar populations at the limits of the PRIMUS and DEEP2 surveys. In orange, we show the CSI selection limit of $3.6\\mu$m $=21$ mag (solid) and the effective optical limit of $i=24.5$ mag (dashed). \\label{fig:bz}} \\end{figure}  While a detailed analysis of this evolution is beyond the scope of this data paper, the power of CSI's selection depth and field size is illustrated by dissecting the sample in some basic ways. Figure \\ref{fig:UVJ} presents CSI bicolor diagrams for several ranges of stellar mass and redshift. We have used our incompleteness functions and the $1/V_{max}$ method to compute the space densities of galaxies, represented using the grayscales. By fixing the grayscale ranges to that derived in the low redshift bin, one can directly see the evolution in number density at fixed mass that was shown in Figure \\ref{fig:phiz}.  These $UVJ$ diagrams are directly comparable to those by other authors \\citep[e.g.][]{williams2009,whitaker2011}, with the red clump of quiescent galaxies and the blue star forming sequence seen previously. The star forming sequence is especially prominent at low stellar masses. Using our H$\\alpha$ and [OII]3727\\AA\\ emission line luminosities to estimate rates of on-going star formation \\citep[e.g.,][after correcting for dust extinction]{kennicutt1998,kewley2004}, we identify those galaxies with high levels of star formation activity ($SSFR>2\\times 10^{-10}$ yr$^{-1}$ at $z=0.1$, evolving to $5.5\\times 10^{-10}$ yr$^{-1}$ at $z=1$) and show those with blue contours. These diagnostics reconfirm that the star-forming sequence in the $UVJ$ diagram remains an efficient means of selecting star-forming galaxies.   The red polygon shows the locus of galaxies with $-1\\le [Z/H]\\le 0.3$ and $A_V\\le 0.5$ and $SSFR < 10^{-11}$ yr$^{-1}$ at $z=0.1$ ($2.8\\times 10^{-10}$ yr$^{-1}$ at $z=1$). We can use this region to isolate galaxies with little, or no on-going star formation (such loci can be generated for such regions in any color space for any filter set). At late times there appears to be a tail out of the high-mass quiescent clump to blue colors along tracks such as those in Figure 4 of \\cite{patel2011} for old populations with very small recent star formation activity. But as Figure \\ref{fig:phiz} already showed, there is a significant diminishing of galaxies with stellar masses $M>M*/2$ at high redshift, and because quiescent galaxies dominate at those mass scales, the red clump, in particular, is quite degraded by $z\\sim 1$.  Being able to track this evolution as a function of mass requires significant depth and we highlight this point in Figure \\ref{fig:bz}, showing the distribution of rest frame $B$-band absolute magnitude as a function of redshift. The irregular redshift bins used here are defined to have constant volume over the CSI sample, such that all bins have similar levels of cosmic variance \\citep[$\\sim 15\\%$;][]{somerville2004}. The CSI $3.6\\mu$m flux limit is shown by the solid orange lines, along with our effective optical limit of $i\\approx 24.5$ mag as dashed orange lines. This $i$-band limit, where our spectroscopic completeness falls below 1/3 of the maximum (of 50\\% at $i=21$ mag, see Fig. \\ref{fig:comp}), dominates over the IRAC selection at $z>1.2$. For comparison, we plot the DEEP2 and PRIMUS magnitude limits of $r=24.1$ mag and $i=23.0$ mag in Figure \\ref{fig:bz} using the blue and green lines, respectively. The red lines in Figure \\ref{fig:bz} show the 10th, 50th, and 90th percentiles of the \\emph{stellar mass density} residing in quiescent galaxies. By $z=1$, DEEP2 or PRIMUS are already missing half the stellar mass in quiescent galaxies.   Surveys like DEEP2 or PRIMUS probe unbiased samples only down to the present day stellar mass of the Milky Way at $z=0.9$, and half that at $z=0.75$. Although CSI's additional reach to $z=1$ for these latter galaxies Way may only extend 1 Gyr farther back in cosmic time, the comoving volume covered within a constant area on the sky is nearly double that of PRIMUS. Such gains in sensitivity highlight the importance of an infrared selection in crafting an unbiased picture of galaxies since $z=1.5$. Furthermore, a methodology for measuring redshifts that does not depend crucially on the presence of emission lines at faint magnitudes, ensures that the resulting samples are relatively free from bias against galaxies with low $M/L$ ratios at fixed mass.   To better illustrate this point, Figure \\ref{fig:mz} shows a grayscale representation of the number densities of galaxies at a given stellar mass and redshift. The 10th, 50th, and 90th percentiles of the stellar mass density residing in quiescent galaxies are shown once again by the red lines. We model the selection differentially as a function of redshift, using the restframe colors of galaxies in each redshift slice to determine which ones get excluded by the survey flux limits at progressively higher redshifts (described in detail in a forthcoming CSI paper on the evolution of stellar mass and star formation), and derive linear fits to the evolution of these percentiles of the quiescent galaxy population. This evolution, or lack therefore, is shown using the dashed black lines in Figure \\ref{fig:mz} and is also visualized in the color-mass diagrams of Figure \\ref{fig:ugm}. The space densities of galaxies at a given restframe $u-g$ color and stellar mass are shown in three redshift bins in CSI, with orange hatched regions marking the projection of the CSI flux limits. The red contours trace the space densities of the quiescent galaxy population (see Figure \\ref{fig:UVJ}). The 10th, 50th, and 90th percentiles in the stellar mass density for quiescent galaxies are overlaid using the vertical dashed lines, with the $0.2<z<0.4$ reference shown by the dotted lines. Our selection in the near-IR, and depth in the optical, ensures that CSI can trace at least 90\\% of the mass in quiescent galaxies back to early times. Each of these panels is normalized individually (unlike Figure \\ref{fig:UVJ}), such that the grayscales and contour levels reflect consistent levels relative to the number densities at their respective epoch. The lack of any significant evolution in these percentiles (or in the red contours) indicates that those galaxies that become quiescent do so without preference for a particular mass scale (at least for galaxies above $M\\gg 10^{10}M_\\odot$), even as the number density of quiescent galaxies has more than doubled over the past 8.5 Gyr.   \\begin{figure}[h] \\centerline{ \\includegraphics[width=2.75in]{f24.pdf} } \\caption{The distribution of our $3.6\\mu$m-selected galaxies in stellar mass, as functions of redshift The green and blue lines indicate the approximate loci for old stellar populations at the limits of the PRIMUS and DEEP2 surveys. In orange, we show the CSI selection limit of $3.6\\mu$m $=21$ mag (solid) and the effective optical limit of $i=24.5$ mag (dashed). The red contours trace the 10th, 50th, and 90th percentiles of stellar mass density residing in quiescent galaxies. By carefully modeling the selection limits, we have computed corrections to these percentiles as functions of redshift and have estimated their linear evolution with redshift. The resulting, underlying evolution of the 10th, 50th, and 90th percentiles of the quiescent galaxy mass density are shown by the dashed black lines. Our selection in the near-IR, depth in the optical, observing strategy, reductions, and analysis of the data enable CSI to reach more than a magnitude fainter in the optical than DEEP2 or PRIMUS for passively evolving galaxies at $z=1$, enabling us to accurately trace the bulk of the stellar mass of passive and active galaxies to $z\\sim 1.4$ in an unbiased fashion. \\label{fig:mz}} \\end{figure}    \\begin{figure*}[t] \\centerline{ \\includegraphics[width=6.5in]{f25.pdf} } \\caption{Restframe color vs stellar mass for galaxies in three redshift bins in CSI, with the grayscale representing the number densities of galaxies at a given color and stellar mass. The 10th, 50th, and 90th percentiles in stellar mass density for quiescent galaxies at each epoch are shown by the vertical dashed lines. The vertical dotted lines show the equivalent percentiles derived from the $0.2<z<0.4$ subsample, highlighting the very low level of evolution in the overall shape of the quiescent galaxies mass function since $z\\sim 1$. The orange hatched regions show the range of stellar mass and color where the CSI flux limits impinge on the sample. The red contours show (arbitrary) levels of galaxy number density for galaxies selected to be quiescent based on the $UVJ$ polygons of Figure \\ref{fig:UVJ}. CSI is tracing the majority of stellar mass on the red sequence, with minimal bias against passive galaxies back almost 9 Gyr in cosmic time. \\label{fig:ugm}} \\end{figure*}      We have described the methodology behind, and initial data from, an ambitious new survey of approximately $2\\times 10^5$ galaxies over a volume that encompasses the last 9 Gyr of cosmic history.  By selecting a sample at 3.6$\\mu$m from Spitzer-IRAC, the CSI survey is the most uniform to-date in terms of limiting mass as a function of redshift, and will ultimately cover an unbiased volume comparable to SDSS but at $0.5<z<1.5$.  The power of the \\emph{IMACS} spectrograph with its low-resolution prisms allows us to survey large volumes efficiently and with sufficient spectral resolution to detect large-scale structure and to measure emission lines from strongly star-forming galaxies and AGN.  CSI provides comparable redshift accuracies for red sequence and blue cloud galaxies; a significant advantage compared to many broadband photometric---and even some spectroscopic---studies.  By combining our low-resolution spectrophotometry with extended broad-band photometry and sophisticated SED modeling, CSI bridges the gap between surveys that are deep enough to probe galaxies below M$^*$ but are small in sample size and volume, and those large-sample, large-volume surveys that do not reach typical, Milky Way-like galaxies that are the main event in the history of cosmic evolution. In forthcoming papers we will update with more detail what we have learned from the first batch of CSI survey data.   One of the highest priority goals will be to take advantage of this large volume to simultaneously measure the evolution of the group mass function \\citep[e.g.][]{williams2012}, and the dependence of star formation on environment \\citep[e.g.][]{kelson2014}. In doing so CSI will uncover the extent to which the growth in the number density of passive galaxies over the past two-thirds of the lifetime of the universe can be attributed to group-related processes such as, for example, galaxy-galaxy interactions (leading to differential number density evolution between group and field galaxies), and/or a decrease in the available fuel for on-going star formation (leaving galaxy number densities unchanged). Given that most galaxies in virialized or massive groups (or both) are ``red and dead'' even at $z\\sim 1$ \\citep{kelson2014}, the mass selection, high spectral coverage, and color-insensitive redshift accuracy of CSI will all be critical elements in characterizing relatively poor groups at high redshift. Subsequent CSI analyses will also focus on, for example, AGN and any connections to group and/or host galaxy properties, on differences between centrals and satellites \\citep[e.g.][]{berlind2002,weinmann2006,vdb2007}, and on merger rates \\citep[e.g.][]{williams2011} as functions of environment."
    },
    "1402/1402.4154_arXiv.txt": {
        "abstract": "We have obtained spectra of \\nqso\\ quasars at $\\mzem>4.4$ with the Gemini Multi Object Spectrometers on the Gemini North and South telescopes, the largest publicly available sample of high-quality, low-resolution spectra at these redshifts. From this homogeneous data set, we generated stacked quasar spectra in three redshift intervals at $z \\sim 5$. We have modelled the flux below the rest-frame Lyman limit ($\\lambda_\\mathrm{r} < 912$\\,\\AA) to assess the mean free path \\lmfp\\ of the intergalactic medium to \\ion{H}{1}-ionizing radiation. At mean redshifts $z_\\mathrm{q} = 4.56, 4.86$ and $5.16$, we measure $\\mlmfp =\\left(\\amfp, \\bmfp, \\cmfp\\right)h_{70}^{-1}$ proper Mpc with uncertainties dominated by sample variance. Combining our results with \\lmfp\\ measurements from lower redshifts, the data are well modelled by a simple power-law $\\mlmfp = A\\left[\\left(1+z\\right)/5\\right]^\\eta$ with $A=\\left(\\aval\\right)h_{70}^{-1}$\\,Mpc and $\\eta = \\gval$ between $z=2.3$ and $z=5.5$. This rapid evolution requires a physical mechanism -- beyond cosmological expansion -- which reduces the cosmic effective Lyman limit opacity. We speculate that the majority of \\ion{H}{1} Lyman limit opacity manifests in gas outside galactic dark matter haloes, tracing large-scale structures (e.g.\\ filaments) whose average density (and consequently neutral fraction) decreases with cosmic time. Our measurements of the strongly redshift-dependent mean free path shortly after the completion of \\ion{H}{1} reionization serve as a valuable boundary condition for numerical models thereof. Having measured $\\mlmfp \\approx 10$\\,Mpc at $z=5.2$, we confirm that the intergalactic medium is highly ionized by that epoch and that the redshift evolution of the mean free path does not show a break that would indicate a recent end to \\ion{H}{1} reionization. ",
        "introduction": "\\label{sec:intro}  The current cosmological paradigm posits that $\\approx 1$\\,Gyr after the Big Bang compact sources -- stars, accreting black holes -- generate sufficient ionizing radiation to reionize the neutral hydrogen throughout the bulk of the intergalactic medium (IGM). Indeed, the spectra of distant quasars and gamma-ray bursts reveal a forest of \\lya\\ absorption lines which are characteristic of a highly ionized medium at $z \\la 6$ \\citep[e.g.][]{gp65,cpb+05}. Resolving the epoch of reionization, its timing and the nature of these ionizing sources stands as one of the outstanding challenges of modern cosmology.  While the community eagerly awaits results of low-frequency radio observations to probe the reionization epoch via the 21cm hyperfine transition \\citep[e.g.][]{zaroubi12,yatawatta13,pober13,beardsley13}, researchers have been studying effects of reionization on the IGM in absorption through spectroscopy of background sources. These include the most distant quasars \\citep{Fan06,mortlock11} and $z>8$ gamma-ray bursts `revealed' by their extremely faint optical fluxes \\citep{kka+06,cucchiara11a,chornock13}. Analysis of these data suggest the transition to a predominantly ionized universe occurs at $z \\ga 6$ \\citep{wbf+03,bolton11}, where a sharp increase in the effective \\lya\\ opacity may occur (but see \\citet{brs07}).  Preferably, one would trace evolution in the ionization state of the IGM in the Lyman continuum which has an effective opacity nearly four orders of magnitude smaller than \\lya\\ and is therefore far more sensitive to the neutral fraction of hydrogen. The effective Lyman Limit (LL) opacity is frequently represented by the mean free path \\lmfp, defined here to be the physical distance a packet of ionizing photons can travel before encountering an $\\mathrm{e}^{-1}$ attenuation. By definition, \\lmfp\\ approaches zero as one transitions into the reionization epoch \\citep[e.g.][]{gnedin00}. After reionization, the mean free path is set by the distribution and evolution of residual neutral gas in the universe.  This will include dense, collapsed structures (e.g.\\ galaxies) but the opacity may be dominated by more diffuse gas in the outskirts of dark matter haloes \\citep[e.g.][]{fumagalli11a} and even more distant and diffuse structures in the IGM. By assessing the redshift evolution of \\lmfp, one constrains the nature of structures providing the universe's LL opacity and, as importantly, its attenuation of the ionizing sources which generate the extragalactic UV background \\citep{flz+09,hm12}.  Many previous works have estimated \\lmfp\\ by first evaluating the \\ion{H}{1} frequency distribution $f(\\mnhi)$ of IGM absorbers (and its redshift evolution; \\fnz) and then convolving this distribution with the photoionization cross-section \\citep[e.g.][]{mm93}. Varying \\fnz\\ within the observational uncertainties, this indirect approach has yielded estimates of 40 to 150\\hmpc\\ (proper) at $z \\sim 3$ in a $\\Lambda$CDM cosmology \\citep{mhr99,flh+08,mm93}. This approach is affected by uncertainty in \\fnz\\ from line blending and clustering of absorbers contributing to the LL opacity, at least at $z\\sim 2.5$ \\citep{rudie13,pmof14}. Recently, the mean free path has been \\textit{directly} evaluated through the analysis of stacked rest-frame quasar spectra: at $z \\approx 4$ using data taken by the Sloan Digital Sky Survey \\citep[SDSS;][]{sdssdr7} and at $z \\sim 2-3$ with space and ground-based programmmes \\citep{pwo09,omeara13,fop+13}. These results show a monotonic increase in \\lmfp\\ with decreasing redshift, ranging from $\\approx 30$\\,\\hmpc\\ at $z=4$ to over 200\\,\\hmpc\\ at $z=2.5$.  For $z>4$, the current constraints on \\lmfp\\ are much poorer owing to the small sample of quasars observed at sufficient signal-to-noise (S/N) to assess the Lyman continuum opacity. Our own analysis of the SDSS data set terminated at $z=4.2$ \\citep{pwo09}, and the SDSS-III survey offers few new bright sources at these redshifts \\citep{paris13}. \\citet{songaila10} surveyed 25 quasars for Lyman limit systems (LLSs) with medium resolution Keck spectroscopy and discovered 20 absorbers with $\\mtll \\ge 1$ at $z>4$, but sampled only 10 systems at $z>4.5$. Combining their results with previous surveys \\citep{storrie94,peroux_dla03}, they measured the incidence of LLSs to $z \\sim 5$ and inferred a mean free path of 50\\hmpc\\ at $z=3.5$ assuming a power-law frequency distribution $f(\\mnhi) \\propto \\mnhi^{-\\beta}$ with a relatively flat $\\beta = 1.3$. Again, \\fnz\\ is poorly constrained at $z>3.5$, especially at column densities $\\mnhi \\approx 10^{17} \\cm{-2}$ implying a significant ($\\sim 30$ per cent) systematic uncertainty in \\lmfp\\ estimates.  In 2010, our group began a multi-semester campaign with the Gemini Multi Object Spectrometers \\citep[GMOSs;][]{gmos} on the twin Gemini 8\\,m telescopes, to survey approximately $150$ quasars at $\\mzem>4.4$, discovered primarily by SDSS. The primary goal of this Giant Gemini GMOS (GGG) survey is to precisely measure \\lmfp\\ at $z\\sim 5$; this is the focus of this manuscript.  The data also enable estimates of the average IGM Lyman series opacity, studies of high-$z$ quasar emission properties, and a search for high-$z$ high-column density absorbers. Those topics will be considered in future manuscripts.  In the following, we adopt a flat $\\Lambda$CDM cosmology with $H_0 = 70 h_{70}\\,\\mathrm{km}\\,\\mathrm{s}^{-1}\\,\\mathrm{Mpc}^{-1}$, $\\Omega_\\mathrm{m} = 0.30$, and $\\Omega_\\Lambda = 0.70$. Unless noted otherwise, all distances quoted in this paper are proper and corrected to the used cosmology. ",
        "conclusions": "\\label{sec:conclusions}  We have presented first results from the Giant Gemini GMOS (GGG) survey of the intergalactic medium at $z\\ga 4$. In a long-term multi-partner Gemini programme we have obtained high-quality (S/N$ \\sim 20$ per $1.85$\\,\\AA\\ pixel) low-resolution (FWHM$\\approx 320\\,\\mkms$) spectra of \\nqso\\ $\\mzem>4.4$ quasars, the largest sample of its kind at these redshifts. The reduced data are publicly available\\footnote{Link to VizieR will be provided upon acceptance of the manuscript.}. The primary goal of this survey has been a precise measurement of the mean free path to \\ion{H}{1} Lyman limit photons in the high-redshift universe via the analysis of stacked rest-frame quasar spectra, a technique pioneered by our team \\citep{pwo09}. Future papers in this series will determine the incidence rate of DLAs and LLSs at $z>4$ and, by using the total Lyman limit opacity determined from the mean free path, constrain the column density distribution of $\\mtll<1$ absorbers. Our primary results are as follows: \\begin{enumerate} \\item The mean free path \\lmfp\\ monotonically decreases with redshift. Subsampling our data set we measure $\\mlmfp=\\left(\\amfp, \\bmfp, \\cmfp\\right)h_{70}^{-1}$ proper Mpc (68 per cent confidence level) at mean redshifts of the stacks $z_\\mathrm{q}=\\left(4.56, 4.86, 5.16\\right)$, respectively (Figs.~\\ref{fig:stacks} \\& \\ref{fig:boot}). \\item When combining these measurements with lower-redshift results obtained with the stacking technique \\citep{pwo09,omeara13,fop+13}, we find that the mean free path smoothly evolves between $z=2.3$ and $z=5.5$ (Fig.~\\ref{fig:mfpvsz}), well described by a power-law: $\\mlmfp\\left(z\\right)=A\\left[\\left(1+z\\right)/5\\right]^\\eta$ with $A=\\left(\\aval\\right)h_{70}^{-1}$\\,Mpc and $\\eta=\\gval$ (68 per cent confidence level). \\end{enumerate}  The redshift evolution of the mean free path exceeds that expected from cosmological expansion, indicating a reduction in number and/or physical size of the absorbing structures with cosmic time. We conclude that a significant fraction of the structures giving rise to \\ion{H}{1} Lyman limit absorption are in the IGM. These structures are likely cosmologically expanding, leading to a substantial decrease of their \\ion{H}{1} fraction with time in the otherwise highly reionized IGM. Although our measurements are consistent with previous, more uncertain estimates based on the statistics of Lyman limit systems and (semi)-numerical models (Fig.~\\ref{fig:mfpfits}), our inferred redshift evolution of the mean free path is very steep, partly due to our correct accounting for cosmological expansion at low redshifts. The smoothly evolving mean free path tracks the Lyman limit absorption in the highly ionized IGM at $2.3<z<5.5$, without any obvious indication of a more rapid decrease at the highest redshifts that would signal the approach of the \\ion{H}{1} reionization epoch. Viable numerical models of \\ion{H}{1} reionization must nevertheless match the measured post-reionization mean free path and its evolution with redshift \\citep[e.g.][]{gnedin06}.  Future work on the mean free path will likely focus on the lowest and highest redshifts. At $z\\sim 1.6$ the mean free path is expected to exceed the horizon of the universe, such that all ionizing sources are expected to contribute to the UV radiation field at any given point \\citep[e.g.][]{mhr99,omeara13}. Currently, there are efforts to constrain this `breakthrough' epoch with \\textit{HST} UV spectroscopy of a sample of $\\mzem\\sim 1$ quasars (Howk et al.\\ in prep.). At the highest redshifts $z\\ga 6$ the mean free path is expected to drop rapidly, indicating the epoch of \\ion{H}{1} reionization \\citep[e.g.][]{gnedin00}. However, measurements of the mean free path are very challenging at $z\\ga 5.5$ due to highly uncertain IGM absorber statistics \\citep{songaila10} and limitations of the quasar stacking technique (high sky subtraction accuracy required for $\\mlmfp\\la 10$\\,Mpc, uncertainty in the UV SED of quasars, bias due to quasar proximity zones). A direct measurement of the mean free path at $z>5.5$ from spectral stacks will require a modest sample of high-quality spectra of either fainter quasars with smaller proximity zones \\citep[e.g.][]{willott10} or the brightest $z\\sim 6$ galaxies \\citep[e.g.][]{willott13} to be collected with future 30\\,m telescopes."
    },
    "1402/1402.0446.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {We present \\strom and near-infrared (NIR) photometry of the bulge cluster NGC~6528 and its surrounding field in the Baade's Window. $uvby$ images have been collected with EFOSC2 on the New Technology Telescope (NTT, La Silla, ESO). The NIR catalogs are based on $J,K$-band VIRCAM@VISTA (Paranal, ESO) and SOFI@NTT photometry. The matching of the aforementioned data sets with Hubble Space Telescope photometry allowed us to obtain proper-motion-cleaned samples of NGC~6528 and bulge stars. Furthermore, we were able to correct the \\stromc--NIR photometry for differential reddening. The huge color sensitivity of the \\stromc--NIR Color--Magnitude--Diagrams (CMDs) helped us in disentangling age and metallicity effects. The red-giant branch (RGB) of NGC~6528 is well reproduced in all the CMDs by adopting scaled-solar isochrones with solar abundance, i.e. $Z = 0.0198$, or $\\alpha$-enhanced isochrones with the same iron content, i.e. $Z = 0.04$, and an age range of $t$ = 10 - 12 Gyr. The same isochrones well reproduce most of the color spread of the Baade's window RGB. These findings support age estimates present in literature for NGC~6528.  We also performed a new theoretical visual--NIR metallicity calibration based on the \\strom index $m_1$ and on visual--NIR colors for red-giant (RG) stars. Scaled-solar and $\\alpha$-enhanced models have been adopted and we validated the new Metallicity--Index--Color (MIC) relations by applying them to estimate the photometric metal abundance of a sample of field RGs and of a metal-poor (M~92, \\feh $\\sim -2.3$) and a metal-rich (NGC~6624, \\feh $\\sim -0.7$) globular cluster. We applied the calibration to estimate the mean metal abundance of NGC~6528, finding \\feh = \\mh = $-0.04\\pm$0.02, with a mean intrinsic dispersion of $\\sigma =$ 0.27 dex, by averaging the metallicities obtained with the scaled-solar $[m],\\ y - J$ and $[m],\\ y - K$ MIC relations, and of $-0.11\\pm$0.01, with $\\sigma =$  0.27 dex, by using the $m_1,\\ y - J$ and $m_1,\\ y - K$ relations. These findings would support the results of Zoccali et al. (2004) based on high-resolution spectroscopy, which give \\feh = $-0.10\\pm$0.2 for NGC~6528, and a low $\\alpha$-enhancement of \\afe = 0.1, and of Carretta et al. (2001), that find  \\feh = 0.07$\\pm$0.01, with a modest $\\alpha$-enhancement, \\afe = 0.2. By applying the scaled-solar MIC relations to the sample of Baade's window RGs, we find a metallicity distribution which extend approximately from \\feh $\\sim -1.0$ up to \\feh $\\sim$ 1 dex, with two metallicity peaks at \\feh $\\approx -0.2$ and \\feh $\\approx$ 0.55 ($[m],\\ y - J$  and $[m],\\ y - K$ relations), and \\feh $\\approx -0.25$ and \\feh $\\approx$ 0.4 ($m_1,\\ y - J$ and $m_1,\\ y - K$ relations). These findings are in fairly good agreement with the spectroscopic studies of Zoccali et al.\\ (2008), and Hill et al. (2011) for the Baade's window, of Uttenthaler et al. (2012) for a region centered at $(l,b) = (0^{\\circ}, -10^{\\circ})$, and with the recent results of the ARGOS surveys for the bulge (Freeman et al.\\ 2013, Ness et al.\\ 2013a).} ",
        "introduction": "The study of the Galactic bulge is fundamental to understand the star formation history of the Galaxy. Moreover, the Milky Way bulge is the closest galaxy bulge that can be observed, and constraining its properties is of great help to the study of more distant galaxies. The Galactic bulge is, indeed, dominated by metal-rich old stars ($t > 10$ Gyr, \\citealt{zocc03}), but the presence of a few intermediate-age stars is not excluded \\citep{vanloon03, bensby13}. The metallicity distribution of red-giant (RG) stars in the bulge spans a range from \\feh $\\approx -1.5$ to 0.5, according to high-resolution spectroscopy of $\\sim$ 800 stars \\citep[hereafter ZO08]{hill11, zocc08}. Furthermore, \\citet{oscar} and \\citet{utten} recently showed that bulge stars are $\\alpha$-enhanced when compared with thin disk stars. Moreover, there is evidence that the  $\\alpha$ enhancement decreases at higher metallicities. More recently, the ARGOS (Abundances and Radial velocity Galactic Origins Survey) survey provided medium-resolution spectroscopy for $\\sim$ 25,000 red clump stars distributed in 28 two-degree fields at different Galactic longitudes and for latitudes $b = -5, -7.5, -10^{\\circ}$. % They found that the bulge stellar content is made by two spatially and chemically distinct sub--populations. {\\em a) \\/} A thin more metal-rich sub--sample, peaking at super-solar iron abundance (\\feh $\\approx$ 0.15), that is kinematically colder and closer to the Galactic plane (A component). This component is associated with the thin disc of the Galaxy. {\\em b) \\/}  A thick more metal-poor sub-sample (peaking at \\feh $\\approx -0.25$) that is uniformly distributed across the selected fields. This component is associated with the boxy/peanut bulge \\citep{ness13}. % They also identify a main more metal-poor component at \\feh = $-0.66$, and two other peaks in their metallicity distribution function (MDF), at \\feh $=  -1.16$, and $-1.73$, that appear to be associated with the thick disk and with the inner halo \\citep{freeman13, ness13}. These results are also in agreement with the high-resolution spectroscopy of a sample of 58 micro-lensed bulge dwarfs by \\citet{bensby13}. Their MDF presents three main \"bumps\" at \\feh $\\approx -0.7, -0.25$, and 0.3, which are in agreement with the three main populations identified by the ARGOS survey. % The above MDFs agree quite well with the results by \\citet{hill11} for the Baade's window (BW). They found, by using high-resolution spectra for 219 bulge red clump giants, a bi-modal metallicity distribution, with peaks at \\feh $\\sim -0.3$ and 0.3. The two main components seem to show different kinematical properties, with the most metal--poor stars more compatible with an old-spheroid, while the metal-rich component with a bar population. A double peaked MDF was also found by \\citet{utten}, based on medium-resolution spectra of $\\sim$ 400 RGs in a bulge field centered at $(l,b) = (0^{\\circ}, -10^{\\circ})$. The two peaks are located at \\mh $\\sim -0.6$ and $\\sim$ 0.3, suggesting that their metal-poor metallicity peak is shifted towards the more metal-poor regime when compared to the metal-poor peaks found by \\citet{bensby11} and \\citet{hill11}. Moreover, they also found that the metal-rich stars show a smaller velocity dispersion compared to metal-poor stars. On the basis of this evidence they suggest that the metal-rich population, making up $\\approx$ 30\\% of the sample, could belong to the bar, while the metal-poor one to the Galactic bulge.  A quantitative explanation of how the Galactic bulge formed is still missing, and medium- and high-resolution spectroscopic studies are tracing different regions of the bulge. Some evidence points to constrain the Milky Way bulge as a classical bulge, others favors a pseudo-bulge. The fact that the RR Lyrae do not show evidence of a bar is also suggesting that the truly old (horizontal branch stars) bulge population forms a spherical bulge, while the intermediate-age population (red clump stars) form a boxy/peanut bulge \\citep{vasquez, dekany}.  In this context, an interesting region to constrain the bulge nature, is the BW. It is characterized by a low-reddening, $E(B-V) <$ 0.6 mag, and by the presence of two Galactic globular clusters (GGCs), namely NGC~6522 and NGC~6528. The former is classified as a bulge metal-intermediate cluster (\\feh $\\sim -1.3/-1.2$ dex). \\citet{io11} estimated a mean metal-abundance of \\feh $= -1.0$ for this cluster, by adopting a \\strom theoretical metallicity calibration based on the $hk$ index. The latter --NGC~6528-- is among the most metal-rich GGCs. Recent spectroscopic investigations, based on high-resolution spectra, suggested for this cluster a solar metallicity and a modest $\\alpha$-enhancement. \\citet{carretta01}, by using four red horizontal-branch (RHB) stars, found \\feh = 0.07$\\pm$0.01 and an $\\alpha$-enhancement of \\afe $\\approx$ 0.2, while \\citet{zocc04}, by using three RGB stars found \\feh = $\u00d00.1\\pm$0.2 and \\afe = 0.1$\\pm$ 0.1. More recently \\citet{origlia}, using high-resolution NIR spectra of four bright giants, found \\feh =  $\u00d00.17\\pm$0.01 and a higher $\\alpha$-enhancement, \\afe = 0.33$\\pm$0.01. The quoted measurements indicate that NGC~6528 is an ideal laboratory not only to constrain the $\\alpha$-enhancement in old metal-rich systems, but also to shed new light on the possible occurrence of an age--metallicity relation among the most metal-rich Galactic globulars (Rakos \\& Schombert 2005; Dotter et al.\\ 2011).  The metallicity estimates, based on photometric indices, of field and cluster stars in the bulge are partially hampered by the occurrence of differential reddening across the field of view and by the mix of stellar populations belonging to the bulge, the thick and thin disk and to globulars. The first problem can be partially overcome by using NIR bands, while the latter is more complex. The use of optical--NIR color--color planes, in particular \\stromc--NIR planes, to separate cluster and field stars \\citep{io09, bono10} is limited by the fact that the metallicity distributions of bulge, thick/thin disk and NGC~6528 all peak around solar chemical composition. Furthermore, the use of the color--color plane to separate field and cluster stars does require precise and deep photometry in at least three optical--NIR bands.  In the current paper we adopt ground--based photometry and use the \\stromc--NIR color--color planes to perform a preliminary selection of bulge, disk and cluster stars. We also take advantage of optical and NIR space photometry collected with Advanced Camera for Surveys (ACS) and Wide Field Camera 3 (WFC3) on board the Hubble Space Telescope (HST) to estimate proper motions and further select the different stellar components. The ground--based \\stromc--NIR photometry is then adopted to estimate the metallicity distribution of bulge and cluster RGs, by applying a new theoretical calibration of the \\strom $m_1 = (v-b) - (b-y)$ metallicity index, based on visual--NIR colors.  This is the first paper of a series devoted to \\stromc--NIR photometry of metal--rich GGCs. In this initial investigation, we derive and test a new theoretical calibration of the $m_1$ index, based on visual--NIR colors for RG stars. The new calibration is adopted to constrain the metallicity distribution of NGC~6528 and its surrounding field.  We favored the adoption of the visual--NIR colors to derive Metallicity--Index--Color (MIC) relations, since they have two clear advantages when compared with indices only based on \\strom colors. {\\em a) \\/} They are not hampered by the presence of $CN$,$CH$, and $NH$ molecular absorption bands. Two strong $CN$ bands ($\\lambda=4142$ and $\\lambda=4215$ \\AA) affects the  $v$ filter, while the strong $CH$-band at $\\lambda=4300$ \\AA~ might affect both the $v$ and the $b$, and the $NH$-band at $\\lambda=3360$ \\AA~ plus the two $CN$ bands at $\\lambda=3590$ and $\\lambda=3883$ \\AA~ might affect the $u$; {\\em ib) \\/} They have a stronger sensitivity to effective temperature. The quoted molecular bands still affect the the new visual--NIR MIC relations, but only through the \\strom $m_1$ index.  The new MIC relations are based on scaled-solar and $\\alpha$-enhanced evolutionary models. The adopted evolutionary models were transformed into the observational plane by adopting bolometric corrections  and color--temperature transformations based on atmosphere models constructed by adopting the same chemical mixtures. They are valid in the metallicity range -2.5 $ < \\feh <$ 0.5. The new MIC relations were validated by estimating the abundances of GGCs, from the very metal-poor M~92 (\\feh $\\sim -2.3$) to the bulge metal-rich cluster NGC~6624 (\\feh $\\sim -0.6$), and of field RG stars, and by comparing them with medium- and high-resolution spectroscopic measurements.  The structure of the current paper is as follows. In \\S 2 we discuss in detail the observations and data reduction, while in \\S 3 we describe the photometric calibration. In \\S 4 we present the \\stromc--NIR catalogs adopted in this paper. Section 5 deals with the approach adopted to calibrate the visual--NIR MIC relations for giant stars, while in \\S 6 we present the different tests we performed to validate the current theoretical MIC relations together with the comparison between photometric estimates and spectroscopic measurements of metal abundances. In \\S 7 we study the metallicity distribution of NGC~6528 and the Baade's window RGs. The summary of the results and a brief discussion concerning further developments of the new MIC relations are given in \\S 8. %_________________________________________________________________________________________ ",
        "conclusions": "We presented \\stromc--NIR photometry of the the bulge globular cluster NGC~6528 and its surrounding field in the Baade's window. The main findings concerning \\stromc--NIR photometry of cluster and field stars are the following:  {\\em a\\/}) The isochrone fit of the proper-motion-cleaned and differential reddening corrected \\stromc--NIR CMDs for NGC~6528 suggests an age of 11$\\pm$1 Gyr, by adopting a scaled-solar isochrone with solar abundance, i.e. $Z = 0.0198$, $Y = 0.273$, or a $\\alpha$-enhanced isochrone with the same iron content, i.e. $Z = 0.04$, $Y = 0.303$. The same scaled-solar and $\\alpha$-enhanced isochrones agree quite well, within the uncertainties, with the NIR CMD of NGC~6528. This result is in good agreement with literature age estimates, i.e. $t =$ 11$\\pm$2 Gyr \\citep{feltz02}, 12.6 Gyr \\citep{momany03}, and 11$\\pm$1 Gyr (LA14).  {\\em b)} We adopted the same theoretical framework used for NGC~6528 to fit \\stromc--NIR CMDs of proper-motion selected and differential reddening corrected bulge stars. We find that a fraction of observed stars along the MS and the RGB are systematically redder and/or fainter than predicted by adopted cluster isochrones. This result would suggest the possible occurrence of large samples of super metal--rich stars, \\feh $>$ 0.1, in the Baade's window. However, these objects could also be explained with an increase either in differential reddening or in depth or both.  {\\em c)} We find that central helium burning stars display a well defined double peak distribution. The fainter peak is located at K$\\sim$13.25$\\pm$0.02 and appears to be associated with the RHB stars, since the peak in the $K$-band luminosity function of cluster RGs in located at K$\\sim$13.16$\\pm$0.02 mag. The brighter (K$\\sim$12.86$\\pm$0.02) peak seems to be associated with intermediate-age RC stars. The two peaks also show a clear separation in color (RHB, J-K$\\sim$1.0; RC, J-K$\\sim$0.9). More quantitative constraints concerning the age and the metallicity distribution of the above evolutionary features do require detailed synthetic CMDs and luminosity functions (Calamida et al., in preparation).  Furthermore, we also provided a new theoretical metallicity calibration based on the $m_1$ index and on visual--NIR colors to estimate the global metal abundance of cluster and field RG stars. We adopted scaled-solar and $\\alpha$-enhanced evolutionary models. This is the first time that visual--NIR colors are adopted to estimate photometric metallicities of RG stars.  We validated the new theoretical metallicity calibration by adopting a sample of field RGs with \\strom and NIR photometry and high-resolution spectroscopy available. The sample includes 61 RGs selected from the study by ATT94 and ATT98. The mean difference between photometric and spectroscopic abundance is $-0.16\\pm0.05$ dex, with a mean intrinsic dispersion of $\\sigma =$ 0.24 dex ($m_1,\\ y-J$, and, $m_1,\\ y-K$ relations), and $0.13\\pm0.06$ dex, with $\\sigma =$0.25 dex ($[m],\\ y-J$ and  $[m],\\ y-K$ relations).  The quoted independent comparisons indicate that the new theoretical MIC relations provide accurate metal abundances for field RG stars with a dispersion smaller than 0.3 dex.  We also tested the calibration by adopting RG stars of two GGCs covering a broad range in metal abundance, i.e. M~92 (\\feh =  $-2.31$) and NGC~6624 (\\feh = $-0.69$), for which both \\strom and NIR photometry were available.  We find that the metallicity distributions of M~92 based on the $\\alpha$-enhanced visual--NIR MIC relations are in very good agreement with the spectroscopic estimates available in literature, i.e. \\feh = $-2.24\\pm$0.10 \\citep{zinn}, and \\feh = $-2.35\\pm$0.05 \\citep{carretta09}. We find, indeed, a global metallicity of  \\mh = $-2.05\\pm$0.04 (\\feh = $-2.40$), with a mean intrinsic dispersion of $\\sigma =$ 0.31 dex, by averaging the values obtained by adopting the $[m],\\ y - J$ and the $[m],\\ y - K$ relations, and of $-2.07\\pm$0.01 (\\feh = $-2.42$), with $\\sigma =$  0.25 dex, by using the $m_1,\\ y - J$ and $m_1,\\ y - K$ relations.  By applying the $\\alpha$-enhanced visual--NIR MIC relations to the bulge cluster NGC~6624 we find averaged mean peaks of $-0.39\\pm$0.09 (\\feh = $-0.74$) and a mean intrinsic dispersion of $\\sigma =$ 0.42 dex ($[m],\\ y - J$, $[m],\\ y - H$, and $[m],\\ y - K$ relations), and of $-0.40\\pm$0.04 (\\feh = $-0.75$), with $\\sigma =$  0.42 dex ($m_1,\\ y - J$, $m_1,\\ y - H$, and $m_1,\\ y - K$ relations). These estimates are in good agreement, within uncertainties, with spectroscopic estimates from the literature (\\feh = $-0.63\\pm$0.09, \\citealt{heasley}, and \\feh = $-0.69\\pm$0.02, \\citealt{valenti11}).  We also apply the new calibration to study the metallicity distribution of the bulge cluster NGC~6528 and its surrounding field. The smoothed metallicity distribution obtained by applying the scaled-solar $[m],\\ y - J$ MIC relation to the proper-motion selected and differential reddening corrected cluster RGs shows a main peak at \\feh = $-0.02$, with a dispersion of $\\sigma =$ 0.19 dex. The metallicity distributions obtained by applying the other scaled-solar MIC relations agree with each other, with average mean peaks of $-0.04\\pm$0.02 and a mean intrinsic dispersion of $\\sigma =$ 0.27 dex ($[m],\\ y - J$ and $[m],\\ y - K$ relations), and of $-0.11\\pm$0.01, with $\\sigma =$  0.27 dex ($m_1,\\ y - J$ and $m_1,\\ y - K$ relations). The agreement between photometric and spectroscopic estimates for the four RGs in common with the spectroscopic samples of ZO04 and SA12 is, within the uncertainties, very good. By adopting the $\\alpha$-enhanced MIC relations to estimate the metallicities of NGC~6528 RGs we obtain metal abundances of more than 0.5 dex too metal-rich, i.e. $\\Delta \\mh = \\mh_{\\hbox{\\footnotesize phot}} - \\mh_{\\hbox{\\footnotesize spec}} \\approx$ 0.5 dex. This result would support the spectroscopic measurements of ZO04 for NGC~6528, which give $\\feh = $-0.10$\\pm$0.2 and a low $\\alpha$-enhancement, i. e. \\afe $\\approx$ 0.1, and the findings of \\citet{carretta01}, based on high-resolution spectroscopy of RHB stars, which give $\\feh =$ 0.07$\\pm$0.01, with a modest $\\alpha$-enhancement, \\afe $\\approx$ 0.2.  The smoothed photometric metallicity distribution of proper-motion selected and differential reddening corrected RGs in the Baade's window surrounding NGC~6528, based on scaled-solar $[m], \\ y - J$ MIC relation, is clearly bimodal, with a main peak at \\feh $\\approx -0.25$ and a secondary one at \\feh $\\approx$ 0.5. Similar distributions are obtained when adopting the other visual--NIR MIC relations, with average peaks at \\feh $\\approx -0.2$ and \\feh $\\approx$ 0.55 ($[m],\\ y - J$, $[m],\\ y - H$ and $m_1,\\ y - K$ relations), and \\feh $\\approx -0.25$ and \\feh $\\approx$ 0.4 ($m_1,\\ y - J$, $m_1,\\ y - H$ , and $m_1,\\ y - K$ relations). By applying the $\\alpha$-enhanced MIC relations we obtain a metallicity distribution shifted by almost 1 dex towards the metal-rich regime.  The shape of our photometric metallicity distribution is quite similar to the shape of the spectroscopic distribution derived by ZO08 and \\citet{hill11} based on a sample of $\\approx$ 200 RGs in the Baade's window. The metal-rich peak of the photometric distribution is slightly shifted ($\\approx$ 0.15 - 0.2 dex) towards the metal-rich regime. The shift could be due to the effect of the molecular $CN/CH/NH$-bands on the $m_1$ index, so the RGs mimic to be more metal-rich than what they really are. Moreover, our MIC relations are valid in the metallicity range $-2.3 <$ \\feh $<$ 0.3, and can be safely adopted in the range $-2.5 <$ \\feh $<$ 0.5, taking into account the current uncertainties in spectroscopic abundances and in the metallicity scale. The metal-rich peak of our distribution is at the edge of the metallicity range covered by our MIC relations.  It is worth noticing that our photometric metallicity distribution for the RGs in the Baade's window is also in fairly good agreement with the spectroscopic MDF obtained by the ARGOS survey for RGs in the Galactic bulge at latitude $b = -5^{\\circ}$ \\citep[see panel a) of their Fig.~11]{ness13}, with the spectroscopic MDF derived by \\citet{bensby13} for a sample of 58 micro-lensed bulge dwarfs, and with the MDF obtained by \\citet{utten} for a sample of $\\sim$ 400 RGs in the a region of the bulge centered at $(l,b) = (0^{\\circ}, -10^{\\circ})$.  The above findings indicate that \\stromc--NIR photometry can provide solid constraints on the metallicity distribution not only of cluster RGs but also of field and halo bulge RGs. The key advantage of the current approach is that it can be applied to a large sample of field RGs located along different line of sight. The main drawback is that it requires accurate and deep photometry in at least three \\strom bands ($v,b,y$) and at least one NIR band. However, accurate NIR photometry is now available for a significant fraction of the Galactic bulge (VISTA, Saito et al.\\ 2012). An extensive use of the current approach is hampered by the fact that current wide field imager at 4--8m class telescopes are not equipped with high-quality \\strom filters. It goes without saying that a medium--large \\strom survey of the low-reddening regions of the Galactic bulge can have a substantial impact on our understanding of the bulge stellar populations. This means more quantitative constraints on the formation and evolution of the bulge and its chemical enrichment."
    },
    "1402/1402.0682_arXiv.txt": {
        "abstract": "Assuming the existence of a primordial asymmetry in the dark sector, we study how  DM-baryon long-range interactions, induced by the kinetic mixing of a new $U(1)$ gauge boson and the photon, affects the evolution of the Sun and in turn the sound speed profile obtained from helioseismology. Thanks to the explicit dependence on the exchanged momenta in the differential cross section (Rutherford-like scattering), we find that dark matter particles with a mass of $\\sim 10\\;{\\rm GeV}$,  kinetic mixing parameter of the order of  $10^{-9}$ and a mediator with a mass smaller than a few MeV  improve  the agreement between the best solar model and the helioseismic data without being excluded by direct detection experiments. In particular, the \\LUX\\ detector will soon be able to either  constrain or confirm our best fit solar model in the presence of a  dark sector with long-range interactions that reconcile helioseismology with thermal neutrino results. ",
        "introduction": "The standard $\\Lambda$CDM cosmological model has been successfully applied to explain the main characteristics of the Universe~\\citep[see e.g.][]{2012arXiv1212.5226H, 2013arXiv1303.5076P}. In particular,  numerical simulations of collisionless cold Dark Matter (DM) describe the gravitational growth of  infinitesimal primordial density perturbations, probed by the cosmic microwave background anisotropies, that lead to the formation of the present-day large-scale structure of the universe. These simulations  provide us with a number of predictions about the structure of cold DM halos and their basic properties~\\citep[see e.g.][]{2013MNRAS.428.1351G}. Although the standard cosmological model has proven to be highly successful in explaining the observed large-scale  structure of the universe, it has been less successful on  smaller scales. Recent data on low mass galaxies suggests that the inferred subhalo DM distributions have almost flat cores~\\citep[see e.g.][]{2010AdAst2010E...5D}, in contradiction with the cuspy profile distributions predicted by numerical simulations~\\citep[see e.g.][]{2010MNRAS.402...21N},  and that sub halo numbers are overpredicted at both low mass and intermediate masses~\\citep[see e.g.][]{1999ApJ...522...82K} and massive dwarf galaxy scales~\\citep[see e.g.][]{2013MNRAS.433.3539G}.  Some of the issues associated with dwarf galaxies can be addressed if the DM is more \"collisional\" (with baryons) than currently considered in numerical simulations. {This hypothesis favors constant-density cores with lower central densities than those coming from cold DM models, such as self interacting DM~\\citep[see e.g.,][]{2013MNRAS.430...81R}.}   In addition, the closeness between $\\Omega_{\\rm DM}$ and $\\Omega_{\\rm b}$, usually referred as cosmic coincidence, may suggest a  profound link between the dark  and the ordinary sectors. Indeed, although the two sectors have different macroscopic properties, the total amount of DM observed today could be  produced  in the early universe by a mechanism identical to baryogenesis, and therefore  an asymmetry between DM particles and their antiparticles would be expected. A detailed account of current progress in asymmetric DM studies can be found in the literature~\\citep[see e.g.,][]{2013arXiv1305.4939P}.    \\medskip All of these cosmological facts may suggest that  dark and ordinary matter may have  more properties in common than expected. In view of this, it is    tantalizing to imagine that the dark world could be similarly complex (CP violating and asymmetric) and full of forces that are invisible to us (hidden parallel sector or mirror world\\footnote{The idea of a mirror world was suggested prior to the advent of the Standard Model (see e.g. Refs.~\\citep{Lee:1956qn,KobzarevOkunPomeranchuk}. The idea that the mirror particles might instead constitute the DM of the Universe was  discussed in refs.~\\citep{Blinnikov:1982eh, 1983SvA....27..371B}.}).  The history of the early mirror universe has been studied in ref.~\\citep{Berezhiani:2000gw}, while the impact of an hidden mirror sector in the CMB and LSS data can be found in ref.~\\citep{Berezhiani:2003wj}. For a general review on the properties of a hidden world neighboring our own, see e.g.~refs.~\\citep{Foot:2004pa, Berezhiani:2005ek}.   More specifically, since in our sector only long--range electromagnetic force and gravity affect the dynamical evolution of  virialized astrophysical objects, the physics of a complex dark sector in which the matter fields are charged under an extra $U(1)$ gauge group is particularly interesting to study. Indeed, if the mass of the new gauge boson (dark photon) is smaller that the typical  momenta exchanged in the scattering, the phenomenology of the dark world  can in itself be as complicated as that of our sector (e.g. dark electromagnetism with quite large self-interactions), providing at the same time also a feeble long-range interaction between the two worlds (thanks to the kinetic mixing of  the new $U(1)$ gauge boson with the photon). This implies in general that the dark world is more collisional than the standard cold DM, and therefore some of the previously mentioned problems can probably be solved~\\citep[see e.g.][]{2011PhRvL.106q1302L}.  In addition, since a long-range  DM-nucleus interaction is  enhanced for small  momentum exchanges, this class of models can also relax the tension between the positive results of direct detection experiments (the annual modulation observed by \\DAMA~\\citep{2008EPJC...56..333B,2010EPJC...67...39B} and \\CoGeNT~\\citep{2011PhRvL.107n1301A,2011PhRvL.106m1301A}, and the hints of an observed excess of events on \\CRESST~\\citep{Angloher:2011uu} and \\CDMSSi~\\citep{2013arXiv1304.4279C,Agnese:2013cvt}), and the constraints coming from null results (e.g.~\\CDMSGe~\\citep{2010Sci...327.1619C}, \\XENON~\\citep{2011PhRvL.107m1302A} and very recently \\LUX~\\citep{Akerib:2013tjd}). The phenomenology of long-range DM-nucleus interactions in direct DM searches has been studied in Refs.~\\citep{2004PhRvD..69c6001F, Foot:2008nw, Foot:2010rj, Foot:2012rk,2011PhRvD..84k5002F, Panci:2012qf}.   \\medskip In this paper,  we investigate how a DM-baryon Long Range Interaction (DMLRI), induced by the kinetic mixing of a new $U(1)$ gauge boson and the photon, affects the evolution of the Sun. The results obtained are then confronted with helioseismology data. The helioseismic data that we use was obtained by several international collaborations,  such as the  Solar and Heliospheric Observatory (SOHO) mission and the  Birmingham Solar Oscillations Network (BiSON) observational network~\\citep{1997SoPh..175..247T,2009ApJ...699.1403B}. Furthermore, we also discuss the constraints on the main parameters of long-range DM particle interactions which can be obtained from direct DM search data. ",
        "conclusions": "We have examined how DM-baryon long--range interactions can affect the Sun's sound speed radial profile. The phenomenological approach used in our analysis lets  us explore all the parameter space from the contact regime to the long-range regime.  We find that DM particles lighter than 8.5 GeV coupled with ordinary baryons through a kinetic mixing parameter $\\gamma_\\phi$ bigger than $5\\times 10^{-9}$, produce a very large impact on the Sun's core in the long--range regime ($m_\\phi$ smaller than few MeV). Therefore they can be excluded as possible DM candidates. However, solar models for which the DM particle has a mass of $10\\;{\\rm GeV}$ and the mediator a mass smaller than $1\\;{\\rm MeV}$ improve the agreement  with helioseismic data. Nevertheless, when the mass of the dark photon is larger than $10\\;{\\rm MeV}$ the impact on the Sun's structure is very small, being in this case the interaction of a contact-type (standard spin-independent picture). In particular, the results obtained here reveal that DM models featuring a long-range interaction with ordinary matter can affect the sound-speed radial profile and in turn probably solve the so-called solar abundance problem without being excluded by terrestrial experiments (e.g. \\LUX\\ and \\XENON) in which the scattering cross section is suppressed by the strong inverse dependence on $q$. The Sun is in fact an ideal  DM detector for the type of particle considered here, since it measures the entire  nuclear recoil energy spectra in  the scatterings.  \\medskip To summarize, in this work we have obtained two main results. Firstly, for the first time, a DM-baryon velocity dependent total cross-section has been implemented  in  solar simulation software. Secondly, but even more importantly, our  analysis shows that DM particles with a mass of 10 GeV and a long--range  interaction with ordinary matter mediated by a very light mediator (below roughly  a few MeV), can have an impact on the Sun's sound speed profile without  violating the constraints coming from direct DM searches. {Our results are valid if the parameter $k_\\chi$ which controls the self-interaction is of the order of $k_\\chi^\\Omega$. Larger values of this parameter, as we have shown in the last paragraph of Sec.~\\ref{Sec:4}, can dramatically reduced the effective luminosity carried by DM particles in the Sun and therefore our results can not be applied in that case.}   Furthermore, as commented upon in Sec.~\\ref{CConstraints}, it might well be possible that a dark sector with long--range forces can  dissipate a relevant amount of energy through the emission of  dark photons.  A dissipative dark sector is extremely interesting because it can offer a rich array of new ideas ranging from the detection of primordial dark radiation to new possibilities for  DM dynamics in virialized astrophysical objects."
    },
    "1402/1402.6292_arXiv.txt": {
        "abstract": "The neutrino signal from SN1987A confirmed the core-collapse scenario and the possible formation of a neutron star. Although this compact object has eluded all observations, theoretical and numerical developments have allowed a glimpse of the fate of it. In particular, a hypercritical accretion model has been proposed to forecast the accretion of $\\sim 0.15 \\:\\mathrm{M_{\\odot}}$ in two hours and the subsequent submergence of the magnetic field in the newborn neutron star. In this paper, we revisit the Chevalier's model in a numerical framework, focusing on the neutrino cooling effect on the supernova fall-back dynamics.  For that, using a customized version of the FLASH code, we carry out numerical simulations of the accretion of matter onto the newborn neutron star in order to estimate the size of the neutrino-sphere,  the emissivity and luminosity of neutrinos.  As a signature of this phase, we estimate the neutrinos expected on SK neutrino experiment and their flavor ratios.  This is academically important because, although currently it was very difficult to detect 1.46 thermal neutrinos and their oscillations, these fingerprints are the only viable and reliable way to confirm the hypercritical phase. Perhaps new techniques for detecting neutrino oscillations arise in the near future allowing us to confirm our estimates. ",
        "introduction": "It took almost 400 years so that another supernova, after Kepler's supernova in 1604, could be observed by the naked eye. The neutrino signal from SN1987A detected in terrestrial observatories (Mont Blanc, Kamiokande, IBM and Baksan), suggested the formation of a neutron star (NS) by the supernova core \\citep{Woosley1988}. So far the search for this holy grail of the type II supernovae paradigm continues. SN1987A was a core-collapse supernova with known progenitor and an estimated distance of $d\\simeq50\\pm1\\:\\mathrm{kpc}$ \\citep{Panagia2005}, in the Large Magellanic Cloud (LMC). The progenitor was a blue supergiant of spectral class B3Iab with 20 $\\mathrm{M_{\\odot}}$, effective temperature $1.6\\times10^{4}\\:\\mathrm{K}$, and an estimated size of 43 $\\mathrm{R_{\\odot}}$. The total energy in the collapse was $\\simeq10^{53}$ erg, but only $1\\%$ of this energy was released in the shock wave of the supernova.\\\\ In the first few seconds, the debris of core-collapse was dense enough that photons could hardly escape.  However, newly created neutrinos in the formation of the newborn NS  escaped carrying out the gravitational energy. The observables of this NS were  an inferred energy $\\sim3\\times10^{53}\\:\\mathrm{ergs}$, temperature $T\\sim4\\:\\mathrm{MeV}$ and decay time-scale of the neutrino burst $\\sim4\\:\\mathrm{s}$. These parameters were consistent with models in which a degenerate iron core collapsed to form a NS \\citep{Burrows1986}, although there is not yet any evidence of the presence of a pulsar, or even a quiet NS inside SN1987A. The possible scenarios to solve this problem range from the delayed collapse of the NS into a black hole \\citep{Brown-Bethe1994}, a demagnetized NS by hyper-accretion of material \\citep{Bernal2013} to the formation of a quark star \\citep{Muslimov-Page1995}.\\\\ In the core-collapse scenario, the shock wave sweeps the outer layers of the progenitor until it encounters a discontinuity in density.  At this point, a reverse shock is generated leading to a fall-back phase which can induce a hypercritical accretion onto the newborn NS surface a few hours after the explosion. This scenario is only possible if the progenitor had a tenuous H/He envelope surrounding a dense He core, as SN1987A \\citep{Smartt2009}.  \\citet{Chevalier1989} argued in favor of such scenario of late accretion onto the SN1987A core and developed an analytical model for the hypercritical regime. In this model, the neutrino cooling plays an important role in the formation of a quasi-hydrostatic envelope around the newborn NS. Recently,  \\cite{Bernal2013, Bernal2010} showed that the magnetic field, for SN1987A parameters, does not play an important role in such regime because it is submerged into the crust of the newborn NS, therefore we will adopt the idea of a demagnetized newborn NS as compact remnant in this supernova. \\\\ Due to the inverse beta decay, electron-positron annihilation ($e^-+e^+\\to Z\\to \\nu_j+\\bar{\\nu}_j$) and nucleon-nucleon bremsstrahlung ($N+N\\to N+N+\\nu_j+\\bar{\\nu}_j$) for $j=e,\\nu,\\tau$, thermal neutrinos will be produced at the core and then they will propagate inside the star. The  properties of these neutrinos will get modified when they propagate in a magnetized medium and depending on the flavor of the neutrino, would  feel a different effective potential because electron neutrino ($\\nu_e$) interacts with electrons via both neutral and charged currents (CC), whereas muon ($\\nu_\\mu$) and tau ($\\nu_\\tau)$ neutrinos interact only via the neutral current (NC). This would induce a coherent effect in which a maximal conversion of $\\nu_e$ into $\\nu_\\mu$ ($\\nu_\\tau$) takes place even for a small intrinsic mixing angle. The resonant conversion of neutrino from one flavor to another due to the medium effect is well known as the Mikheyev-Smirnov-Wolfenstein effect \\citep{wol78}. These  neutrino oscillations have been widely studied in the literature in different scenarios \\citep{1999A&A...344..573R,1987ApJ...314L...7G,2000APh....13...21V,2012ARNPS..62..407J,2013MNRAS.434.1355J}.\\\\ In this work we do a numerical study of hyper-accretion of matter onto the newborn NS surface in the SN1987A scenario. We will consider as input the analytical model proposed by \\citet{Chevalier1989} and \\citet{Brown-Weingartner1994} about the late accretion of matter onto compact objects, few moments after the core-collapse, focusing on the importance of the neutrino cooling processes on the NS surface. We employ the AMR FLASH code to carry out the numerical simulations of the reverse shock and the complex dynamics near stellar surface, including several neutrino cooling processes, a more detailed equation of state, and an additional degree of freedom in the system. As a signature of this phase we estimate the neutrino luminosity,  the number of events as well as the flavor ratio that could have reached the Earth. The paper is arranged as follows. In section \\ref{section2} we describe the analytical model as input on the numerical approach focusing on the neutrino cooling processes present in the system. In section \\ref{section3} we develop the neutrino oscillations models. In section \\ref{section4} we discuss the radial profiles inside the supernova remnant and the neutrino propagation in such regions. Finally, In section \\ref{section5} we present and discuss the results in the SN1987A framework. ",
        "conclusions": "\\label{section5} We have studied hypercritical accretion onto newborn NS with the aim of looking into the neutrino cooling effect in the formation of an envelope in quasi-hydrostatic equilibrium of SN1987A. This is particularly relevant in the context of making the NS eventually invisible as a pulsar following the supernova explosion, this with the condition that  the magnetic field is submerged by the accretion in the hypercritical regime.  Extending our exploration of parameter space to spherical symmetry spanning a significant fraction of the stellar surface in the simulations,  we have focused on the size of the neutrino-sphere where the emissivity of neutrinos is effective and then calculated the neutrino luminosity in the first instant after the reverse shock has reached the hard surface of the star and the formation of the quasi-steady atmosphere takes place.\\\\ We made comparisons of our numerical results with the analytical model of Chevalier and found an excellent agreement between the neutrino luminosity and height scale values obtained with the Flash Code and those estimated with analytical approximations. The additional neutrino processes involved in the numerical approach slightly increase the neutrino luminosity value and the height scale. This is because the analytical model of the hypercritical system was performed with some assumptions while the numerical approach has included more physical ingredients (magnetic field, a more realistic equation of state and additional neutrino processes).  We also analyze other interesting phenomena such as the submergence of the magnetic field on the stellar surface as well as the dynamics of the system with higher degrees of freedom.\\\\ On the other hand, taking into account the obtained neutrino luminosity and parameters of the SN-II, we estimated that 1.49 neutrinos must have reached the detector two hours later after the main neutrino burst.  However, because this number is much smaller that the $\\sim 20$ events could have been undetected.   On the other hand, if considering the value of volume 31.9 kton in  SK-III  \\citep{2012PhRvD..86a2006R}, we would expect a neutrino number of 22.3 in the hypercritical phase. The amount of these neutrinos would confirm such phase in other SN with the identical characteristics (temperature, distance, etc) We have studied the active-active neutrino process in the supernova framework. We have divided the path of moving neutrinos into four regions.  For the first region, we have used the neutrino effective potential derived in \\citet{2014arXiv1401.1581F} which is  a function of chemical potential ($\\mu$), temperature ($T$),  neutrino energy ($E_\\nu$)  and magnetic field.   We have shown that for a neutrino test of energy 1, 5 , 20 and 30 MeV, and parameters considered of temperature and chemical potential in the range of 1 MeV $\\leq T \\leq$ 5 MeV,   $10^{-2}\\, {\\rm eV} \\leq \\mu  \\leq 10^3\\, {\\rm eV}$,  neutrinos oscillate resonantly, for two- and three-neutrino mixing.  In regions from two to four, we have also calculated each effective potential and then analyzed their oscillations through each zone.   For three neutrino mixing,  we have calculated the ratio flavor expected on Earth. Our analysis shows that deviations  from 1:1:1 are obtained  for neutrino energies of 5, 10, 15 and 20 MeV, given in table 1.  Distinct flavor ratios  will provide constraints on parameters of this hypercritical accretion.  Although currently it is very difficult to detect  neutrino oscillations, new techniques for detecting neutrino oscillations arise in the near future allowing us to confirm our estimates and thus the hypercritical phase.\\\\ It is important to highlight that we estimate a detectable number of neutrinos assuming an object at the distance of SN 1987A at the LMC, however it could be used for future predictions in a more favorable situation where a core collapse might take place in the Milky Way Galaxy."
    },
    "1402/1402.4859_arXiv.txt": {
        "abstract": "We show that rapidly-spinning black holes can display turbulent gravitational behavior which is mediated by a new type of parametric instability. This instability transfers energy from higher temporal and azimuthal spatial frequencies to lower frequencies--- a phenomenon reminiscent of the inverse energy cascade displayed by $2+1$-dimensional turbulent fluids. Our finding reveals a path towards gravitational turbulence for rapidly-spinning black holes, and provides the first evidence for gravitational turbulence in an asymptotically flat spacetime. Interestingly, this finding predicts observable gravitational wave signatures from such phenomena in black hole binaries with high spins and gives a gravitational description of turbulence relevant to the fluid-gravity duality. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.6637_arXiv.txt": {
        "abstract": "{ Super-Earths with solid surfaces such CoRot-7b and Kepler-10b are expected to be extremely hot. It has been suggested that they must be atmosphere-free and that a lava ocean is present on their hot dayside. Here, we use several dedicated thermal models to study how observations with NIRSPEC on the JWST could further confirm and constrain, or reject the atmosphere-free lava ocean planet model for very hot super-Earths. Using CoRoT-7b as a working case, we explore the consequences on the phase-curve of a non tidal-locked rotation, with/without an atmosphere, and for different values of the albedo. We simulate future observations of the reflected light and thermal emission from CoRoT-7b with NIRSPEC-JWST and look for detectable signatures, such as time lag, as well as the possibility to retrieve the latitudinal surface temperature distribution. We demonstrate that we should be able to constrain several parameters after observations of two orbits (42 h) with a broad range of wavelengths: i)The Bond albedo is retrieved to within +/- 0.03 in most cases. ii) The lag effect allows us to retrieve the rotation period within 3 hours for a planet, with rotation half the orbital period. iii) Any spin period shorter than a limit in the range 30 - 800 h, depending on the thickness of the thermal layer in the soil, would be detected. iv) The presence of a thick gray atmosphere with a pressure of one bar, and a specific opacity higher than $10^{-5} m^{-2} kg^{-1}$ is detectable. v) With spectra up to 4.5 $\\mu$m, the latitudinal temperature profile can be retrieved to within 30 K with a risk of a totally wrong solution in 5 \\% of the cases. We conclude that it should thus be possible to distinguish the reference situation from other cases. Surface temperature map and the albedo brings important constraints on the nature or the physical state of the soil of hot super-Earths. } ",
        "introduction": "During the last decade, the development of high precision photometry, particularly from space, led to one of the most exciting, albeit expected, accomplishment:  the first detection of massive rocky planets, or the so-called super-Earths (SE). Despite their large masses, the terrestrial nature of these super-Earths has been confirmed by the measure of their mean densities by combining both transit and radial velocity observations. Surprisingly, the first two objects that are observed compatible with a rocky composition, CoRoT-7b \\citep{Leger_2009, Queloz_2009, Hatzes_2011}  and KEPLER-10b \\citep{Batalha_2011}, appear to be similar in radius ($1.68 \\pm 0.09$ and $1.42 \\pm 0.03  \\, \\mathrm{R_\\oplus}$), mass ($7.42 \\pm 1.21$  and $4.56 \\pm 1.2  \\, \\mathrm{M_\\oplus}$), and environment (They both orbit around a main sequence G star with a very short period roughly equal to 0.85 day). However, their formation is far from being understood. In particular, whether they could have formed {\\it in situ} or should have migrated to this close orbit or whether they formed as rocky planets or as gas giants whose gaseous envelope subsequently evaporated \\citep{Leitzinger_2011} remains unclear.  Considering the magnitude of the energy dissipation induced by the unusually strong tides, this kind of rocky objects should most likely be phase-locked in a very short timescale ($<1$ My). Furthermore, the extreme UV fluxes and stellar winds received by the planet ensure an efficient atmospheric escape. As a result, if an atmosphere is present, it is probably very thin and in condensation/sublimation equilibrium with the molten surface, although the maximum possible pressure of such an atmosphere is still debated (\\citealt{SF09}; \\citealt{Leger_2011} -hereafter L11-; \\citealt{CM11}; \\citealt{Kite_2011}).  According to L11, \\citet{Mura_2011}, and \\citet{Rouan_2011}, the direct expected consequences on the planet surface properties are as follow:  \\begin{itemize} \\item A violent thermal gradient between the day/night hemispheres (2500-3000 K versus 50 K) because of the absence of thermal surface redistribution. \\item A large lava ocean on the day side, where the temperature exceeds the rock fusion temperature. \\item A cometary-like tail trailing the planet \\citep{Mura_2011}. These tails have already been detected for Mercury (e.g. \\citealt{Potter_2002}) and possibly for HD 209458 b  \\citep{Vidal-Madjar_2003}. Depending on the atmospheric molecular composition, this tail can be observable. Constraints on the molecules present in the atmosphere of exoplanets have already been inferred from observations with the VLT \\citep{Guenther_2011}. \\end{itemize}  To get a better understanding of the extreme nature of this new class of objects, we need more observations. In particular, estimations of the albedo and of the surface fraction occupied by the lava ocean could constrain the composition and the structure of the crust of the planet. The challenging point lies in the difficulty of the direct observation of the emitted/reflected light from these objects because of their very small radius and of the extremely high contrast with respect to the star they orbit. However, \\citet{Rouan_2011} and \\citet{Samuel_2011} have shown, respectively, on the example of KEPLER-10b and CoRoT-7b, that the modulation of the planetary emitted flux due to the phase effect should be detectable in a near future with the James Web Space Telescope (JWST). With that in mind, we decided to explore the variety of phase curves that an observer based in the Earth neighborhood would see during the orbit of the considered SE in different physical situations. We decided to focus on the case of CoRoT-7b as a typical example of a very hot SE around a sun-like star. One could easily extend this study with the same method to any other very hot SE, while adapting the properties of the system as the star and the planetary radia, their masses, the magnitude of the star, the orbital period, and so on. ",
        "conclusions": "In the context of an increasing number of detected very hot super-Earths, such as Corot-7b or Kepler-10b,  we have explored whether future observations with the instrument NIRSpec on JWST would allow us to constrain their surface properties.  In particular, we have studied how the presence of a dense atmosphere or of a non synchronous rotation of the planet could be inferred from multiwavelength light curves. We also studied to which accuracy one could retrieve the value of the surface Bond albedo and the temperature distribution using the SED between 0.8 and 3.5 $\\mu m$.  We developed a model of the thermal behavior of the planet surface by assuming that a lava lake is maintained around the substellar point and by considering  various cases of heat penetration and albedo. This is done to estimate the time lag of the thermal emission when the planet is not phase-locked and observed all along its orbit. We also modeled the effect of a putative atmosphere at various pressures on the surface temperature using an GCM  specifically developed for the study of extrasolar planets. We  finally examined how the actual distribution of temperature along a meridian could be retrieved from the observed SED by assuming a polynomial distribution with latitude.  Thanks to the broad range of wavelengths accessible with JWST and provided that  a signal to noise ratio around 5 per resel is reached after a total exposure time of $ \\sim$ 70 h,  we are generally able to constrain several of those parameters with a fair accuracy : \\begin{itemize} \\item The Bond albedo is retrieved to within $\\pm 0.03$ in most cases. \\item Thanks to the asymmetry of the light curve due to the lag effect, the rotation period of a non phase-locked planet is retrieved with an accuracy of 3 hours   when the rotation is fast, or half the orbital period. For a longer period, the accuracy is reduced. \\item The shortest period compatible with observation of a phase-locked planet is in the range 30 - 800 h depending on the thickness of the heat penetration layer. This means that any rotation period shorter than this limit would be detectable. \\item We should be able to detect  the presence of a gray atmosphere, provided it is thick enough: with a pressure of one bar and an specific opacity higher than $10^{-5} m^{-2} kg^{-1}$. \\item The latitudinal temperature profile can be retrieved with an accuracy of 10 to 50 K depending on the signal to noise ratio obtained with NIRSPEC (20 to 5 assumed). We note that the risk of a totally wrong solution is not excluded but is at maximum of 5\\% for a SNR of 5 and a wavelength range extending to 4.5 $\\mu m$. \\end{itemize} Note that the four first results  would still be valid if the exposure time was to be reduced to two complete orbital periods (i.e., 42h).  These results show that we could test the most current hypothesis of phase-locking and of an atmosphere-free planet for planets similar to Corot-7b and confirm or reject the lava ocean model of L11 by using NIRSPEC-JWST spectra. A precise determination  of the albedo and temperature distribution of the planet surface appears also  doable, giving the  possibility to constrain more precise models of the planet structure and properties. The possible discovery of very hot exolanets around  stars  brighter than Corot-7 or Kepler-10 would offer better opportunities to test those predictions: for instance, a more affordable integration time of only several hours would  be needed for a V = 8 solar-type star.  Of course, one cannot exclude that our model is too simplistic and/or that the effects we did not envision could change some of the conclusions on the feasibility.  For instance, the hypothesis concerning a constant albedo, which is independent of wavelength, the nature and the phase (lava or solid) of the soil or, even more so, the hypothesis concerning a gray gas on the atmosphere may seriously affect the amplitude of the time lag effect we predict if the actual situation is significantly different. On the other hand, we consider the two results on the mean abedo derivation and on the possibility to  retrieve the latitudinal temperature profile to a large extent as rather solid.  The access to the phase curves simultaneously in visible and near IR seems to be a promising source of information that will provide useful constraints on the very hot super-Earths surfaces and in general on hot exoplanets."
    },
    "1402/1402.3823_arXiv.txt": {
        "abstract": "{} {We present r-Java 2.0, a nucleosynthesis code for open use that performs r-process calculations as well as a suite of other analysis tools. } {Equipped with a straightforward graphical user interface, r-Java 2.0 is capable of; simulating nuclear statistical equilibrium (NSE), calculating r-process abundances for a wide range of input parameters and astrophysical environments, computing the mass fragmentation from neutron-induced fission as well as the study of individual nucleosynthesis processes.} {In this paper we discuss enhancements made to this version of r-Java, paramount of which is the ability to solve the full reaction network.  The sophisticated fission methodology incorporated into r-Java 2.0 which includes three fission channels (beta-delayed, neutron-induced and spontaneous fission) as well as computation of the mass fragmentation is compared to the upper limit on mass fission approximation.  The effects of including beta-delayed neutron emission on r-process yield is studied.  The role of coulomb interactions in NSE abundances is shown to be significant, supporting previous findings.  A comparative analysis was undertaken during the development of r-Java 2.0 whereby we reproduced the results found in literature from three other r-process codes.  This code is capable of simulating the physical environment of; the high-entropy wind around a proto-neutron star, the ejecta from a neutron star merger or the relativistic ejecta from a quark nova.  As well the users of r-Java 2.0 are given the freedom to define a custom environment.  This software provides an even platform for comparison of different proposed r-process sites and is available for download from the website of the Quark-Nova Project: \\url{http://quarknova.ucalgary.ca/}   } {} ",
        "introduction": "\\label{intro} A key nucleosynthesis mechanism for the production of heavy elements beyond the iron peak is rapid neutron capture, or the r-process \\citep{burb57,cam57}.  In spite of a large volume of observational data \\citep{sneden08}, the astrophysical site(s) of the r-process remains an open question.  Explosive, neutron-rich environments provide the ideal conditions for r-process to occur.  The predominant astrophysical sites which are being studied as possible locations for r-process are: the high-entropy winds (HEW) from proto-neutron stars \\citep[see][]{qianWoos96,farouqi10} and ejected matter from neutron star mergers \\citep[see][]{fried99,gori11}.  However both of these scenarios face significant hurdles that must be overcome.  The long time-scale for neutron star merger events \\citep{faber12} limits this scenario's ability to explain the r-process element enrichment of metal-poor stars \\citep{sned03}.  HEW models have been shown to be very sensitive to the chosen physical conditions of the winds and elaborate hydrodynamic models have yet to prove that the HEW scenario can provide the necessary environment for significant r-process to occur \\citep[e.g.,][]{hoff08,janka08,rob10,fisc10,wanajo11}.  The theoretical quark nova has as well been proposed as a potential r-process site \\citep{jaik07}.  The explosive, neutron-rich environment of a quark nova provides an intriguing avenue of study for astrophysical r-process \\citep{jaik07}.  It is important to note that the observed nuclear abundance of r-process elements in metal-poor stars \\citep{sned03} and abundance data of certain radionuclides found in meteorites \\citep{qian08} point to the likelihood of multiple r-process sites \\citep{tru02}.  Certainly there remains much to be learned about astrophysical r-process and its necessary conditions.  To help drive this study we have developed r-Java \\citep{rjava1}, which is a cross-platform, flexible r-process code that is transparent and freely available for download by any interested party\\footnote{r-Java 1.0 \\& 2.0 can be downloaded from quarknova.ucalgary.ca}.  The purpose of this article is to introduce the second version of r-Java (r-Java 2.0), discuss the new features which it contains and display a selection of simulation results.  Prior to delving into the details of r-Java 2.0 it would be enlightening to briefly examine the capabilities and limitations of the first version of the code.  As discussed in \\cite{rjava1} our aim for r-Java was to create an easy-to-use, cross-platform r-process code that avoids the $^{\\backprime}$black-box$^{\\prime}$ pitfalls that plague many scientific codes.  In order to achieve this goal, r-Java was developed with an intuitive graphical user interface (GUI) and we provide extensive documentation as well as user tutorials on our website. The first version of r-Java was predicated upon the waiting point approximation (WPA) which assumes an equilibrium between neutron captures and photo-dissociations.  The practical effect of using the WPA is that the relative abundance along isotopic chains depends only on neutron density ($n_n$), temperature ($T$) and neutron separation energy ($S_n$), \\citep[For more details see][eqn. 12]{rjava1}.  Through the WPA the number of coupled differential equations that must be solved is reduced from thousands to on the order of a hundred (the number of isotopic chains in the reaction network), greatly reducing computational costs.  The utilization of the WPA comes at the expense of generality, in that the assumption is only valid in high temperature and neutron density environments (typically only considered for $T_9 > 2$ ,where $T_9$ is in units of $10^9$ K and $n_n > 10^{20}$ cm$^{-3}$).  Within the context of the WPA r-Java 1.0 provided users the ability to run r-process simulations for a wide variety of scenarios, including neutron irradiation of static targets as well as the dynamic expansion of r-process sites.  A key feature of r-Java that has been maintained through the release of the second version is the ability for users to easily make changes to nuclear inputs between simulation runs.  Other flexibilities of r-Java include the ability to specify the amount of heating from neutrinos, turn on/off various processes and choose the density profile of the expanding material.  This paper is organized as follows.  Section \\ref{code} provides an overview of the advancements made in r-Java 2.0 and discusses the default rates included with the code.  The nuclear statistical equilibrium (NSE) module is discussed in section \\ref{nse}.  The new fission methodology is detailed in section \\ref{fiss}.  The effect of $\\beta$-delayed neutron emission is discusssed in section \\ref{bdn}.  Section \\ref{full} compares the full reaction network calculation with the WPA approach.  Comparative analysis between r-Java 2.0 and other full reaction network codes is carried out in section \\ref{bench}.  Finally in section \\ref{summary} a summary is provided along with a look to future work on r-Java. ",
        "conclusions": "\\label{summary}  This paper has discussed the nuclear physics incorporated in r-Java 2.0; providing cutting-edge fission calculations, $\\beta$-delayed neutron emission of up to three neutrons and neutron capture and photo-dissociation rates from one of the most sophisticated mass models (HFB21).  Nevertheless it is the flexibility to change any parameter quickly and easily that makes r-Java 2.0 a powerful tool for the study of nuclear astrophysics.  As well r-Java 2.0 is capable of solving a full r-process reaction network containing over 8000 nuclei and can do so both accurately and efficiently, with a typical full reaction network simulation completed on the order of minutes.  The scientific aim of this release of r-Java is to study r-process nucleosynthesis in the expansion phase ($T \\lesssim 3 \\times 10^9$ K \\cite[e.g.][]{howard93}) and NSE at high temperature ($T \\gtrsim 4 \\times 10^9$ K \\cite[e.g.][]{tru66}).  We are currently developing a charged-particle reaction network module that will be incorporated in a future version of r-Java.  With a more realistic treatment of fission we have added to r-Java 2.0 the ability to investigate the role of fission recycling in the r-process.  In the past by simply using the mass cut-off approach the mistake of going to too high of a neutron density was masked by the fact that fission recycling would not allow the r-process to proceed beyond the cut-off.  This manifests in the r-process abundance at neutron freeze-out in two ways; an under-production of super-heavy nuclei and the over-production of nuclei around the A=130 magic number.  With the fission methodology implemented here the super-heavy regime (A $>$ 270) can be studied using r-Java 2.0. The preliminary study undertaken here supports the findings of \\cite{peter12}, where super-heavy nuclei (A$\\sim 290$) can be formed by the r-process.  The super-heavies subsequently decay on the order of seconds.  The emission of $\\beta$-delayed neutrons can act to maintain a sufficiently high neutron density to allow for the r-process to reach heavier elements.  The effect of $\\beta$-delayed neutron emission is as well significant during the decay to stability once the r-process has stopped.  They act to smooth out the nuclei distribution on the path to stability and shifts the abundances to lower masses.  Their role may in some cases not be as direct as just stated.  The $\\beta$-delayed neutrons can alter the r-process path, accessing nuclei which would more readily capture neutrons causing the neutron density to drop more rapidly than if they were ignored.  This must be studied in more detail and with r-Java 2.0 the user can quickly and easily investigate the effect of $\\beta$-delayed neutrons on r-process abundances.  By performing a comparative study between r-Java 2.0 and three other full network r-process codes we have found good agreement between the codes, however undertaking this analysis has highlighted the potential pitfalls of comparing the results from different codes.  Factors such as; choice of mass model, evolution methodology of physical parameters, code stopping criteria and precision can contribute to variations in r-process abundances that are artefacts of the nucleosynthesis code structure rather than physical scenarios being studied.  This comparative analysis highlights the universality of r-Java 2.0, which by allowing the user to customize both the nuclear and astrophysical parameters, is capable of reproducing the results of other nucleosynthesis codes.  The development of r-Java 2.0 was done in a way to maximize the flexibility of the software, allowing for the adjustment of any nuclear or physical property both quickly and easily.  The choice of Java as the programming language allowed for the inclusion of an easy to use GUI that is cross-platform compatible.  Beyond its applicability to scientific study, the goal of r-Java 2.0 was to make it accessible in a teaching capacity by ensuring it is easy to use and allowing for the investigation of individual processes.  In the follow-up paper to this work we will turn our attention to the astrophysical side of r-process which is well covered by r-Java 2.0.  Built into the interface of r-Java 2.0 is the option to define a custom density evolution or to select one of three proposed astrophysical r-process sites; high-entropy winds around proto-neutron stars \\citep[example studies;][]{woos92,qianWoos96,thomp01,farouqi10}, ejecta from neutron star mergers \\citep[][and others]{freid99b,gori11} or ejecta from quark novae \\citep{jaik07}.  For each of the proposed astrophysical sites r-Java 2.0 consistently calculates the temperature and density evolution, the details of which will be discussed in the aforementioned upcoming paper.  By including the physics of different astrophysical sites in one piece of r-process software we have provided an even platform to compare the r-process abundances of different astrophysical sites."
    },
    "1402/1402.6401_arXiv.txt": {
        "abstract": "We report the results from our analysis of the \\textit{Fermi} Large Area Telescope data of the \\textit{Fermi} unassociated source 2FGL J1906.5$+$0720, which is a high-ranked candidate pulsar. In order to better study our target, we first update the ephemeris for PSR J1907$+$0602, which is used to help remove any possible contamination due to strong emission from this nearby pulsar. From our analysis, 2FGL J1906.5$+$0720 is confirmed to have a significant low energy cutoff at $\\sim$ 1~GeV in its emission (14$\\sigma$--18$\\sigma$ significance), consistent with those seen in young pulsars.  We search for pulsations but no spin frequency signals are found in a frequency range of 0.5--32 Hz. No single model can fully describe the source's overall \\fermi\\ $\\gamma$-ray spectrum, and the reason for this is the excess emission detected at energies of $\\geq$4 GeV. The high-energy component possibly indicates emission from a pulsar wind nebula, when considering 2FGL J1906.5$+$0720 as a young pulsar. We conclude that 2FGL J1906.5$+$0720 is likely a pulsar based on the emission properties we have obtained, and observations at other energies are needed in order to confirm its pulsar nature. ",
        "introduction": "Since the \\textit{Fermi} Gamma-ray Space Telescope was launched in June 2008, the main instrument on-board---the Large Area Telescope (LAT) has been continuously scanning the whole sky every three hours in the energy range from 20 MeV to 300 GeV, discovering and monitoring $\\gamma$-ray sources with much improved spatial resolution and sensitivity comparing to former $\\gamma$-ray telescopes \\citep{atw2009}. In 2012 resulting from \\textit{Fermi}/LAT data of the first two-year survey, a catalog of 1873 $\\gamma$-ray sources was released by \\citet{nol2012} as the \\textit{Fermi}/LAT second source catalog. Among the $\\gamma$-ray sources, approximately 800 and 250 were found to be respectively associated with blazars and active galaxies of uncertain types, and more than 100 were associated with pulsars in our Galaxy. The three types thus account for the majority of the $\\gamma$-ray sources detected by \\textit{Fermi}. In addition, 575 sources in the catalog have not been associated with any known astrophysical objects \\citep{nol2012}. For the purpose of identifying the nature of these unassociated sources, many follow-up studies, such as classifying their $\\gamma$-ray characteristics \\citep{ack2012}, searching for radio pulsars \\citep{ray2012}, and observing at multi-wavelengths \\citep{tak2012,ace2013}, have been carried out.  Because of the relative lack of sources at low Galactic latitudes in many extragalactic source catalogs and the emission contamination by the Galaxy, the Galactic distribution of the \\textit{Fermi} unassociated sources were found to concentrate towards the Galactic plane \\citep{nol2012}. More than half of the unassociated sources are located at low latitudes with $|b|<$ 10\\textdegree \\ \\citep{nol2012}, possibly suggesting Galactic origins for most of them. Taking under consideration the types of identified and associated Galactic $\\gamma$-ray sources in the catalog, these low-latitude unassociated sources are most likely pulsars, pulsar wind nebulae, supernova remnants, globular clusters, or high-mass binaries.  Additionally since identified and associated AGNs or blazars have a nearly isotropic distribution, AGN/blazar origins for these sources can not be excluded. In any case, the low-latitude \\textit{Fermi} unassociated sources are the best young pulsar candidates on the basis of currently known Galactic $\\gamma$-ray populations, as $\\sim$50\\% of the identified or associated Galactic \\fermi\\ sources are pulsars \\citep{nol2012} and the \\textit{Fermi}-detected millisecond pulsars are nearly isotropic (see Figure 2 in \\cite{lat2013}). With high rotational energy loss rates (so-called spin-down luminosities), young pulsars are clustered close to the Galactic plane and can be detected to large distances.  Aiming to search for new pulsars among the unassociated sources, we selected the pulsar candidates from the \\textit{Fermi} second source catalog by requiring $|b|<$ 10\\textdegree \\ and variability indices (Variability\\_Index parameter in the catalog) lower than 41. The variability indices were reported to measure the variability levels of sources, and a value larger than 41.64 indicates $<1$\\% chance of being a steady source \\citep{nol2012}. We further ranked the candidates by their Signif\\_Curve parameters reported in the catalog, which represent the significance of the fit improvement between curved spectra and power-law spectra, as $\\gamma$-ray pulsars typically have curved spectra with a form of exponentially cutoff power law. The first ten sources from our selection are listed in Table~\\ref{tab:candi}. The first source listed is 2FGL J1704.9$-$4618, which has the highest Signif\\_Curve value of $\\sim$9.97$\\sigma$ but the lowest detection significance value ($\\sim$9$\\sigma$; Signif\\_Avg parameter in the catalog). For a comparison, the second source in our list 2FGL J1906.5$+$0720 has both high Signif\\_Curve ($\\sim$9.85$\\sigma$) and Signif\\_Avg values ($\\sim$24$\\sigma$), and is ranked the first among candidate pulsars by \\citet{lee2012}, who applied a Gaussian-mixture model for the ranking. Among the bright $\\gamma$-ray sources ($>$20$\\sigma$ detection significance), this source is clearly located in the pulsar region in the plane of the curvature significance versus variability index \\citep{rom2012}. We thus carried out detailed study of 2FGL J1906.5$+$0720 by analyzing \\textit{Fermi}/LAT data of the source region, and report our results in this paper.  In addition, 2FGL J1906.5$+$0720 is located close to a very bright $\\gamma$-ray pulsar J1907$+$0602 (Signif\\_Avg $\\sim$ 55$\\sigma$; \\cite{lat2013}). The angular distance between them is approximately 1.3 degrees (see Figure~\\ref{fig:ts}). The pulsar was discovered in the first $\\sim$4 month LAT data, revealed with a spin frequency of $\\sim$9.378 Hz and a spin-down luminosity of $\\sim$2.8$\\times$10$^{36}$ erg s$^{-1}$ \\citep{abd2009}. The pulsar is radio faint, making very difficult to study its timing behavior at radio frequencies \\citep{abd2010}. In order to better study our targeted \\fermi\\ source by removing possible contamination from PSR J1907$+$0602, we performed timing analysis to the LAT data of the pulsar and include our timing results in this paper. ",
        "conclusions": "\\label{sec:dis}  By carrying out phase-connected timing analysis of the nearly 5-year \\fermi\\ $\\gamma$-ray data of PSR J1907$+$0602, we have obtained the timing parameters and updated the $\\gamma$-ray ephemeris for this pulsar. The obtained timing parameters are similar to those given in the second \\fermi\\ catalog of $\\gamma$-ray pulsars \\citep{lat2013}. However the glitch decay time constant is $\\sim$99 days, larger than $\\sim$33 days given in the catalog. This difference is likely due to the longer time span of the data we analyzed (2-year data was analyzed in the second \\fermi\\ catalog of $\\gamma$-ray pulsars; \\cite{lat2013}) and the unstable timing parameters caused by the timing noise. PSR J1907$+$0602 is quite young with a characteristic age of $\\sim$19.5 kyr \\citep{abd2010}. The post-fit rms timing residual was 2.1 ms, resulting from our timing analysis (Table~\\ref{tab:timing}).  We performed different analyses of the \\fermi/LAT data for the unassociated source 2FGL J1906.5$+$0720. Through likelihood analysis with different spectral models, we confirmed that a curved spectrum with a low-energy break or cutoff at $\\sim$1 GeV is clearly preferred to a simple power law. The significances of the curvature ($\\sim\\sqrt{TS}\\sigma$) are approximately 14--16 $\\sigma$ and 16--18 $\\sigma$ for the full data and the offpulse phase interval data, respectively. This feature is characteristic of $\\gamma$-ray pulsars detected by \\textit{Fermi}. On the basis of the \\fermi\\ second pulsar catalog, young $\\gamma$-ray pulsars have 0.6 $<\\Gamma<$2 and 0.4 GeV $< E_{cutoff} <$ 5.9 GeV, and millisecond $\\gamma$-ray pulsars have 0.4 $<\\Gamma <$ 2 and 1.1 GeV $< E_{cutoff} <$ 5.4 GeV (\\cite{lat2013}). If 2FGL J1906.5$+$0720 is a pulsar, its Galactic location and spectral feature suggest that it is probably a young pulsar (see, e.g., \\cite{lat2013}). It should be noted that a log parabola, which better fits the spectra of 2FGL J1906.5$+$0720, is usually used to model the spectra of $\\gamma$-ray binaries \\citep{nol2012}. However, considering the non-detection of any long-period modulations and the low variability of 2FGL J1906.5$+$0720, a $\\gamma$-ray binary is not likely the case for the source.   From our spectral analysis, a high-energy component was found to exist at $\\geq$4 GeV in the emission of 2FGL J1906.5$+$0720. Considering it as a young pulsar, the component likely originates from its pulsar wind nebula (PWN; e.g., \\cite{gs06}). A pulsar wind generates a termination shock by the interaction of high-energy particles contained in it with the ambient medium, at which particles are re-distributed and can radiate ultra-relativistic emission. In the $\\gamma$-ray energy range, the \\fermi\\ second source catalog used 69 known PWNe for the automatic source association, and found that nearly all of them (except three) are associated with young pulsars \\citep{nol2012}. However since $\\gamma$-ray emission from a pulsar often dominates over that from its PWN, the number of PWNe that have been confirmedly detected by \\fermi\\ is limited \\citep{ack2011}. For 2FGL J1906.5$+$0720, our spatial distribution analysis has confirmed the existence of the high-energy component in its spectrum, but the putative PWN would be too small or too faint to be resolved by \\fermi. Further X-ray imaging of the source field is needed in order to detect the PWN and thus help verify the pulsar nature for 2FGL J1906.5$+$0720.  We have not been able to find any pulsed emission signals from the \\fermi\\ data of 2FGL J1906.5$+$0720, which is required to verify the source's pulsar nature. We note that the LAT blind search sensitivity depends on many parameters, such as the accurate position of the source, the source region used for pulsation search, contamination from background diffuse emission and from nearby sources (given that our target is located at the Galactic plane with several identifiable sources nearby). Using the sensitivity estimation method for the blind searches provided by \\citet{dor2011}, the pulsed fraction of 2FGL J1906.5$+$0720 should be $\\gtrsim$0.57 for a detection probability of $>$68\\% (the 1-year detection significance is $\\sim$20$\\sigma$ for the source). \\citet{dor2011} also extracted an all-sky detectability flux map to describe the minimum 0.3 -- 20 GeV photon flux required for the detection of pulsars with pulsed fractions. In the inner Galactic plane the detectability flux should be higher than $\\sim$10$^{-7}$ ph cm$^{-2}$ s$^{-1}$. The 0.3 -- 20 GeV photon flux we obtained for 2FGL J1906.5$+$0720 is $<$10$^{-7.1}$ ph cm$^{-2}$ s$^{-1}$ (derived from spectral parameters listed in Table~\\ref{tab:likeli}), suggesting the difficulty of detecting pulsed emission from the source through blind searches. Considering the radio pulsations from the source have been searched several times but with no detection \\citep{ray2012}, in order to verify its pulsar nature, X-ray observations are needed.  \\bigskip  We thank the referee for valuable suggestions. This research was supported by Shanghai Natural Science Foundation for Youth (13ZR1464400), National Natural Science Foundation of China (11373055), and the Strategic Priority Research Program ``The Emergence of Cosmological Structures\" of the Chinese Academy of Sciences (Grant No. XDB09000000). ZW is a Research Fellow of the One-Hundred-Talents project of Chinese Academy of Sciences.   \\clearpage \\begin{table} \\caption{The first 10 candidate pulsars ranked by Signif\\_Curve}\\label{tab:candi} \\begin{center} \\begin{tabular}{lccccccc} \\hline \\hline Source & Signif\\_Curve ($\\sigma$) & $Gb$ (\\textdegree) & Variability\\_Index & Signif\\_Avg ($\\sigma$) \\\\ \\hline 2FGL J1704.9$-$4618 & 10.0 & $-$3.111 & 21.3 & 9.3 \\\\ 2FGL J1906.5$+$0720 & 9.8 & $-$0.002 & 30.9 & 24.0 \\\\ 2FGL J1819.3$-$1523 & 9.2 & $-$0.072 & 30.0 & 19.3 \\\\ 2FGL J1847.2$-$0236 & 8.5 & $-$0.257 & 31.3 & 13.8 \\\\ 2FGL J1856.2$+$0450c & 8.4 & 1.139 & 18.8 & 12.3 \\\\ 2FGL J1619.0$-$4650 & 8.3 & 2.457 & 22.2 & 10.6 \\\\ 2FGL J2033.6$+$3927 & 8.3 & $-$0.382 & 33.8 & 13.0 \\\\ 2FGL J1045.0$-$5941 & 8.3 & $-$0.639 & 21.5 & 36.1 \\\\ 2FGL J0858.3$-$4333 & 8.1 & 1.428 & 16.5 & 14.1 \\\\ 2FGL J1739.6$-$2726 & 8.1 & 1.906 & 27.6 & 15.2 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\clearpage \\begin{table} \\caption{$\\gamma$-ray ephemeris for PSR J1907$+$0602.}\\label{tab:timing} \\begin{center} \\begin{threeparttable} \\begin{tabular}{lc} \\hline \\hline Parameter & Value\\tnote{*} \\\\ \\hline R.A., $\\alpha$ (J2000.0) & 19:07:54.7343205 \\\\ Decl., $\\delta$ (J2000.0) & 06:02:16.97850 \\\\ Pulse frequency (s$^{-1}$) & 9.3776609432(10) \\\\ Frequency first derivative (s$^{-2}$) & $-$7.62737(7) $\\times$ 10$^{-12}$ \\\\ Frequency second derivative (s$^{-3}$) & 1.95(2) $\\times$ 10$^{-22}$ \\\\ Epoch of frequency (MJD) & 55422.275976 \\\\ Dispersion measure (cm$^{-3}$ pc) & 82.1 \\\\ 1st glitch epoch (MJD) & 55422.155 \\\\ 1st glitch permanent frequency increment (s$^{-1}$) & 4.3466(3) $\\times$ 10$^{-5}$ \\\\ 1st glitch frequency deriv increment (s$^{-2}$) & $-$7.72(2) $\\times$ 10$^{-14}$ \\\\ 1st glitch frequency increment (s$^{-1}$) & 2.10(5) $\\times$ 10$^{-7}$ \\\\ 1st glitch decay time (Days) & 99(4) \\\\ rms timing residual (ms) & 2.1 \\\\ Time system & TDB \\\\ E$_{min}$ & 50 MeV \\\\ Valid range (MJD) & 54683--56497 \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\footnotesize \\item[*] Parameters with no uncertainty reported are fixed to the values given in the second \\textit{Fermi} catalog of $\\gamma$-ray pulsar \\citep{lat2013} except the glitch epoch. \\end{tablenotes} \\end{threeparttable} \\end{center} \\end{table}  \\clearpage \\begin{table} \\caption{Maximum binned likelihood results for 2FGL J1906.5$+$0720}\\label{tab:likeli} \\begin{center} \\begin{threeparttable} \\begin{tabular}{llcccccc} \\hline \\hline Spectral model & Parameters & Full data & Offpulse phase interval data & Offpulse phase interval data\\\\ & & & & (`tail' removed) \\\\ \\hline PowerLaw & $\\Gamma$ & 2.31 $\\pm$ 0.02 & 2.42 $\\pm$ 0.02 & 2.31 $\\pm$ 0.02 \\\\ & $G_{\\gamma}$ (10$^{-11}$ erg cm$^{-2}$ s$^{-1}$) & 15 $\\pm$ 0.5 & 13 $\\pm$ 0.4 & 11 $\\pm$ 0.4 \\\\ & TS$_{PL}$ & 1101 & 2437 & 1595 \\\\ \\hline LogParabola & $\\alpha$ & 2.52 $\\pm$ 0.05 & 2.82 $\\pm$ 0.05 & 2.73 $\\pm$ 0.07 \\\\ & $\\beta$ & 0.35 $\\pm$ 0.03 & 0.37 $\\pm$ 0.03 & 0.51 $\\pm$ 0.04 \\\\ & $E_{b}$\\tnote{*} (GeV) & 1 & 1 & 1 \\\\ & $G_{\\gamma}$ (10$^{-11}$ erg cm$^{-2}$ s$^{-1}$) & 13 $\\pm$ 0.5 & 12 $\\pm$ 0.4 & 9 $\\pm$ 0.3 \\\\ & TS$_{LP}$ & 1388 & 2795 & 1966 \\\\ \\hline PLSuperExpCutoff & $\\Gamma$ & 1.7 $\\pm$ 0.2 & 1.7 $\\pm$ 0.1 & 1.2 $\\pm$ 0.1 \\\\ & $E_{cut}$ (GeV) & 1.7 $\\pm$ 0.8 & 1.2 $\\pm$ 0.2 & 0.8 $\\pm$ 0.1 \\\\ & $G_{\\gamma}$ (10$^{-11}$ erg cm$^{-2}$ s$^{-1}$) & 13 $\\pm$ 5 & 12 $\\pm$ 2 & 9 $\\pm$ 1 \\\\ & TS$_{PL+cutoff}$ & 1297 & 2694 & 1853 \\\\ \\hline PLSuperExpCutoff & $\\Gamma$ & 1.4 $\\pm$ 0.3 & 1.4 $\\pm$ 0.1 & 1.4 $\\pm$ 0.2 \\\\ obtained by fitting & $E_{cut}$ (GeV) & 1.0 $\\pm$ 0.3 & 0.5 $\\pm$ 0.1 & 0.6 $\\pm$ 0.1 \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\footnotesize \\item[*] The break energies are fixed at 1~GeV. \\end{tablenotes} \\end{threeparttable} \\end{center} \\end{table}  \\clearpage \\begin{figure} \\begin{center} \\includegraphics[width=3in]{source-1.eps} \\includegraphics[width=3in]{offpulse-source-1.eps} \\end{center} \\caption{TS maps (0.1--300 GeV) of $\\mathrm{5^{o}\\times5^{o}}$ regions centered at R.A.= 286.647$\\mathrm{^{o}}$, Decl.= 6.6$\\mathrm{^{o}}$ (equinox J2000.0) extracted from the full data ({\\it left} panel) and offpulse phase interval data ({\\it right} panel) of PSR J1907$+$0602. The image scales of the maps are 0\\fdg1 pixel$^{-1}$. Two putative nearby sources are marked by circles.} \\label{fig:ts} \\end{figure}  \\clearpage \\begin{figure} \\begin{center} \\includegraphics[width=5in]{postfit_rms.eps} \\end{center} \\caption{Post-fit timing residuals for PSR J1907$+$0602.} \\label{fig:rms} \\end{figure}  \\clearpage \\begin{figure} \\begin{center} \\includegraphics[width=5in]{profile.ps} \\end{center} \\caption{Folded pulse profile and two-dimensional phaseogram in 32 phase bins obtained for PSR J1907$+$0602. For clarity, two rotations are shown on $X$-axis. The gray scale represents the number of photons in each bin.} \\label{fig:ftp} \\end{figure}  \\clearpage \\begin{figure} \\begin{center} \\includegraphics[width=3in]{n0-1906-20.eps} \\includegraphics[width=3in]{n0-1906-offpulse-20.eps} \\end{center} \\caption{$\\gamma$-ray spectra of 2FGL J1906.5$+$0720 extracted from the full data ({\\it left} panel) and offpulse phase interval data ({\\it right} panel). The exponentially cutoff power laws and the log parabolas obtained from maximum likelihood analysis (see Table~\\ref{tab:likeli}) are displayed as dashed and dotted curves. The solid curves are the exponentially cutoff power laws obtained by fitting the data points below 2 GeV.} \\label{fig:spec} \\end{figure}  \\clearpage \\begin{figure} \\begin{center} \\includegraphics[width=3in]{offpulse-1906-removesource-1.eps} \\includegraphics[width=3in]{n0-1906-offpulse-addsource-20.eps} \\end{center} \\caption{{\\it Left panel}: TS map (0.1--300 GeV) of $\\mathrm{5^{o}\\times5^{o}}$ region centered at R.A.= 286.647$\\mathrm{^{o}}$, Decl.= 6.6$\\mathrm{^{o}}$ (equinox J2000.0) extracted from the offpulse phase interval data of PSR J1907$+$0602 when the two nearby sources (the positions are marked by circles) were removed. The image scale of the map is 0\\fdg1 pixel$^{-1}$. {\\it Right panel}: $\\gamma$-ray spectrum of 2FGL J1906.5$+$0720 extracted from the offpulse phase interval data when the two nearby sources were removed. The dashed and dotted curves represent the exponentially cutoff power law and the log parabola, respectively, obtained from maximum likelihood analysis (see Table~\\ref{tab:likeli}). The solid curve represents the exponentially cutoff power law obtained by fitting the data points below 2 GeV.} \\label{fig:notail} \\end{figure}  \\clearpage"
    },
    "1402/1402.3230_arXiv.txt": {
        "abstract": "The accurate classification of galaxies in large-sample astrophysical databases of galaxy clusters depends sensitively on the ability to distinguish between morphological types, especially at higher redshifts.  This capability can be enhanced through a new statistical measure of association and correlation, called the {\\it distance correlation coefficient}, which has more statistical power to detect associations than does the classical Pearson measure of linear relationships between two variables.  The distance correlation measure offers a more precise alternative to the classical measure since it is capable of detecting nonlinear relationships that may appear in astrophysical applications.  We showed recently that the comparison between the distance and Pearson correlation coefficients can be used effectively to isolate potential outliers in various galaxy datasets, and this comparison has the ability to confirm the level of accuracy associated with the data.  In this work, we elucidate the advantages of distance correlation when applied to large databases.  We illustrate how the distance correlation measure can be used effectively as a tool to confirm nonlinear relationships between various variables in the COMBO-17 database, including the lengths of the major and minor axes, and the alternative redshift distribution.  For these outlier pairs, the distance correlation coefficient is routinely higher than the Pearson coefficient since it is easier to detect nonlinear relationships with distance correlation.  The V-shaped scatter plots of Pearson versus distance correlation coefficients also reveal the patterns with increasing redshift and the contributions of different galaxy types within each redshift range. ",
        "introduction": "The classification of galaxies has been of great interest for decades, and the ability to distinguish between morphological types is pertinent, especially  at higher redshifts.  \\citet{kin96} created ultraviolet to near-infrared spectral energy distributions for local galaxies to establish a relationship between morphological type and spectral energy distribution for the mainstream Hubble classification of elliptical, bulge, lenticular, and Sa-Sc spiral galaxies, as well as starburst galaxies.  The overall lack of emission lines hinders the accurate classification of redshifts for more distant galaxies, except in the case of starburst galaxies with optical spectra \\citep{kin96}.  In addition, redshift surveys typically measure inadequate luminosity functions for spiral, elliptical, and lenticular galaxies because of contamination by dwarf galaxies \\citep{del03}; hence there is a need to more accurately distinguish between the giant and dwarf galaxy types.  These concerns suggest that there remains a need to identify objects that may have been misclassified using traditional methods.  Mart{\\'i}nez-G{\\'o}mez, Richards \\& Richards (2014) recently applied a new statistical measure of association and correlation, called the {\\it distance correlation coefficient}, to the COMBO-17 database.  This was the first application of distance correlation to astrophysical data.   Some advantages of the distance correlation measure are that it can detect nonlinear associations that are undetectable by the classical Pearson correlation coefficient, it is applicable to random variables of any dimension, and it is zero if and only if the variables are independent \\citep{sze09,due14}.  Moreover, \\citet{sze09} and \\citet{sim12} have shown that the distance correlation coefficient has higher statistical power to detect associations than the well known Pearson correlation coefficient of \\citet{pea1895}.  We selected a sample of 15,352 galaxies, with redshifts $0 \\le z < 2$ , from the COMBO-17 database.  The analysis was performed on 33 variables, including  mean redshift, luminosity distance, length of major axis, length of minor axis, position angle, magnitudes, and fluxes at various wavelengths and in different observing runs \\citep{mar14}.  The Pearson correlation coefficient and distance correlation coefficient were then compared for the 528 pairs of variables corresponding to the selected 33 variables, and the results were displayed in scatter plots.  These plots have distinctive horseshoe or V-shapes, which occur whenever multi-dimensional data are mapped into two dimensions \\citep{mar14,dia08}.   They provide a mechanism by which the differences between potential outliers and the remaining data points can be accentuated.  \\begin{table*}[t] \\caption{Galaxy Types and Redshift Ranges} \\centering \\begin{tabular}{llccccc} \\hline\\hline \\M Galaxy  & Kinney et al. (1996)              & Magnitude Range based on \t\t\t\t     & \\multicolumn{4}{c}{Number of Galaxies, N} \\M \\\\ \\cline{4-7}\\M Type     & ~~~~~~~Template            & Wolf et al. (2003a)   \t\t\t          &  ~~$0 \\le z < 0.5$~~  &  ~~$0.5 \\le z < 1$~~   & ~~$1 \\le z < 2$~~ & Total \\\\[0.5ex]\\hline \\M Type 1  & Elliptical, bulge, S0, Sa     & {$B-r > 1.25$}, {$m_{280}-B \\geq 1.1$} &  ~~38 & ~~50 & ~~16 & ~~~104\\\\ Type 2  & Spiral: Sa, Sbc          & {$B-r > 1.25$}, {$m_{280}-B < 1.1$}      &  ~~45 & ~~19 & ~~~4 & ~~~~68 \\\\ Type 3  & Spiral: Sbc - SB6                 & {$0.95 < B-r \\leq 1.25$} \t\t\t\t\t          & ~328 & ~277 & ~109 & ~~~714\\\\ Type 4  & Starburst: SB6 - SB1~~  & {$B-r \\leq 0.95$} \t\t\t\t\t\t\t          & 3254 & 9284 & 1928 & 14466\\\\ [0.5ex]\\hline \\M Total     & Elliptical - Starburst    &  \t\t\t\t\t\t\t\t\t\t\t\t\t    & 3665 & 9630 & 2057 & 15352\\\\ [0.5ex] \\hline \\end{tabular} \\label{table1} \\end{table*}  With the availability of this new statistical tool, we can learn more about the COMBO-17 dataset through a deeper examination of the distance correlation results, specifically about the separate classes of galaxies in the sample.  In this Letter, we extend the results of \\citet{mar14} to illustrate the ways in which distance correlation can be used to explore general patterns in the  database, and to identify pairs of variables that are associated with deviations from nonlinearity.  In Section 2, we review the definition of the distance correlation coefficient, identify a more extensive set of potential outlier pairs, and examine the patterns with galaxy type and redshift.  In Section 3, we provide a summary of results and conclusions. ",
        "conclusions": "In this Letter, we extended the work of \\citet{mar14} to further explore the application of distance correlation to large galaxy databases, determine the level of association between the measured properties of these galaxies, and examine the consistency of the galaxy classifications across various redshift groups.  For the application to the COMBO-17 database, we selected a sample of 15,352 galaxies, with redshifts $0 \\le z < 2$, and we studied the associations between 528 pairs of variables, based on a selection of 33 variables for each galaxy.  The comparison between the Pearson correlation coefficient and the distance correlation coefficient creates a V-shaped scatter plot, which is an effective tool for identifying potential outlier pairs of variables.  Most of the potential outliers identified in Table \\ref{table2} are directly linked to nonlinear relationships between the variables in the outlier pair.   We also used the V-shaped scatter plots to examine the levels of accuracy associated with three redshift groups of starburst (Type 4) galaxies, and found that the tighter scatter plot for starburst galaxies at lower $z$ confirms the higher precision of these measurements relative to those at higher $z$.  We have shown that distance correlation is more effective than the classical Pearson coefficient in determining the nature of the association between pairs of variables, and we recommend that the distance correlation coefficient be used in place of the Pearson coefficient to determine the relationships between variables.  Another impact of our results lies in the ability to examine the quality of associations between pairs of variables that are related to large groups of galaxies; this has been achieved through the V-shaped scatter plots that compare the Pearson and distance correlation coefficients, and which depend on the number of galaxies in the sample. If a given galaxy dataset is contaminated by galaxies that do not belong to that group, then the misclassified galaxies may influence the level of association between certain pairs of variables, since some of these relationships may vary with galaxy type.  Hence, knowledge of the level of association between pairs of variables reflects information about the galaxies from which the data were collected, and we can examine the overall quality of the data for groups of galaxies of different types (elliptical, lenticular, spiral, starburst) and also over a range of redshifts.  In the case of the COMBO-17 database, the Type 2 and Type 3 groups of spiral galaxies are likely to be contaminated substantially by starburst galaxies (de Lapparent 2003); these starburst galaxies also dominate the Type 4 galaxy group in our sample.  Our study revealed further evidence of this contamination by means of the distribution of points in the V-shaped scatter plots for the Type 2 and Type 3 galaxies, and this distribution resembles a superposition of multiple V-shaped patterns (Figures 1 and 2; large right frames).  Similar superpositions of V-shaped patterns were also revealed when we separated the database into the four galaxy types and then plotted the information onto a single graph (Figures \\ref{fig1} - \\ref{fig3}; large left frames).   Given the increasingly large sizes of galaxy databases, our technique can be used to quickly identify potentially contaminated galaxy groups and then a more detailed study can be performed to identify the specific misclassified galaxies."
    },
    "1402/1402.1373_arXiv.txt": {
        "abstract": "This work investigates a particular class of artefacts, or ghost sources, in radio interferometric images. Earlier observations with  (and simulations of) the Westerbork Synthesis Radio Telescope (WSRT) suggested that these were due to calibration with incomplete sky models. A theoretical framework is derived that validates this suggestion, and provides predictions of ghost formation in a two-source scenario. The predictions are found to accurately match the result of simulations, and qualitatively reproduce the ghosts previously seen in observational data. The theory also provides explanations for many previously puzzling features of these artefacts (regular geometry, PSF-like sidelobes, seeming independence on model flux), and shows that the observed phenomenon of flux suppression affecting unmodelled sources is due to the same mechanism. We demonstrate that this ghost formation mechanism is a fundamental feature of calibration, and exhibits a particularly strong and localized signature due to array redundancy. To some extent this mechanism will affect all observations (including those with non-redundant arrays), though in most cases the ghosts remain hidden below the noise or masked by other instrumental artefacts. The implications of such errors on future deep observations are discussed. ",
        "introduction": "In the context of radio interferometry, the term {\\em calibration} refers to the estimation and correction of instrumental errors (which are traditionally taken to also include effects of the troposphere and ionosphere) on the observed visibilities. Current calibration approaches boil down to a joint fit to the observations of a sky model and an instrumental model, such as that provided by the radio interferometer measurement equation \\citep[RIME;][]{ME1,RRIME1}. A typical observing strategy will include intermittent {\\em calibrator scans} of a known calibrator field, for which an accurate prior sky model is available; the obtained instrumental solutions can then be interpolated onto scans of the target field. These can be further refined through a process known as {\\em self-calibration} or {\\em selfcal} \\citep{Cornwell:selfcal}. Selfcal is an iterative approach (the sky model is refined at each iteration) that minimizes the error between predicted visibilities corrupted by the instrumental model (the free parameters) and the observed visibilities in a least squares sense during each iteration. An initial sky model for selfcal can be obtained by imaging visibilities that have been corrected by the interpolated calibrator solutions. Where a reasonable initial sky model for the target field is available, it can even provide the starting point for selfcal without the need for an external calibrator.  Traditional selfcal assumes an instrumental model where all effects are direction-independent. The increased field of view of modern radio interferometers implies that direction dependent effects can no longer be ignored during calibration. Incorporating direction dependent effects into calibration solutions (third-generation calibration, or 3GC) has become a major research field over the past few years \\citep{Intema2009,Smirnov2011,Kazemi2011,Kazemi2013,Wijnholds2009}. \\citet{Veen2004,Rau2009} have conducted good literature reviews on calibration.  It is well-established that calibration can lead to imperfect images, even to the generation of spurious source components, elimination or suppression  of real components, and the deformation of the structure of extended sources \\citep{Linfield1986,Wilkinson1988,Taylor1999,Marti2010,Marti2008}. \\cite{Kazemi2013a} have proposed a novel calibration technique that is meant to minimize the amount of source suppression that occurs if the sky model is incomplete. This approach uses a $t$-distribution to model the residual noise. \\cite{Marti2008} has shown that spurious point sources can form when performing calibration on data containing nothing but white noise.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{fig1-gdb-ghosts-crop.pdf} \\caption{Ghost sources in a 92cm WSRT observation of J1819+3845. The target field is in the lower-image corner of the image, and Cyg A is to the upper left, just outside the image.} \\label{fig:2004ghosts} \\end{figure*}   One of the more striking sightings of calibration artefacts occurred in 2004, in a 92cm Westerbork Synthesis Radio Telescope (WSRT) observation of J1819+3845 by de Bruyn (Fig.~\\ref{fig:2004ghosts}). After self-calibration, the map exhibited a string of ``ghosts'' -- point-source-like objects, mostly of negative flux, arranged along a line linking the brightest object in the field with Cyg A, which was about 20$^\\circ$ away (i.e. in a distant sidelobe of the primary beam, and extremely attenuated by WSRT's extremely low sidelobe response). The pattern was highly peculiar as the ghosts were arranged with some regularity, and their positions did not vary with frequency. No other observations at the time were known to exhibit such features, and the problem remained open until a series of 21cm WSRT observations in 2010, which were done as part of the ``Quality Monitoring Committee'' project \\citep{Smirnov2011qmc}. In these observations, a large pointing error was deliberately introduced, and the resulting residual images (post self-cal) exhibited similar artefacts. The ghosts were fainter, but there were several strings of them, all associated with the brightest objects in the field (Fig.~\\ref{fig:2010ghosts}). The problem was then investigated empirically, through the use of simulations \\citep{Smirnov2010ghosts}, and this revealed a number of features: \\begin{itemize} \\item The ghosts were associated with sky model errors (i.e. missing or incorrect flux in the sky model, or direction-dependent errors towards the brightest sources). In the QMC case, this was due to the large pointing error; in the J1819+3845 case this was due to insufficiently accurate modelling of Cyg A. Correcting for these errors (by solving for differential gains towards the brighter sources in the QMC case, and towards Cyg A in the J1819+3845 case) made the ghosts disappear. \\item A simple simulation of a two-source (1 Jy and 1 mJy) field, where only the 1 Jy source was included in the calibration model, while the second source played the role of ``contaminator'', produced a similar ghost pattern in the residual visibilities. The peak intensity of the pattern was roughly at $\\mu$Jy level, and appeared to be proportional to the flux of the contaminator source (but did not depend on the flux of the model source!) This suggested that ghosts should {\\bf always} arise in the presence of incomplete sky models, but would generally be buried in the thermal noise, unless the observations were very sensitive, or the missing model sources were sufficiently bright. \\item The ghosts always arranged themselves along a line (or lines) passing through the unmodelled or poorly modelled source(s), and the dominant source(s) in the sky model. The positions of the ghosts corresponded to some (but not all) rational fractions of the interval between the sources (i.e. 1/2, 1/3, 2/3, 1/5, etc.), with significant variation in intensity. The positions did not depend on frequency. \\item The ghosts exhibited sidelobes that were similar, but not identical to, the PSF of the telescope. \\item Similar simulations with other telescopes (VLA) showed a far less regular artefact pattern. \\end{itemize}  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{fig2-qmc-ghosts-crop.pdf}  \\caption{Ghost sources in a 21cm WSRT observation of the QMC2 field. Note that this is a residual dirty map, i.e. the sky model sources have been subtracted, and the visibly brighter sources here are in fact relatively faint. The positions of the brightest subtracted sources are indicated by red circles. Multiple strings of ghosts are visible, here highlighted by the ellipses. Note how the strings are firmly associated with the brightest sources.} \\label{fig:2010ghosts} \\end{figure*}  These facts strongly suggested that the regularity of the ghost patterns in Figs.~\\ref{fig:2004ghosts} and \\ref{fig:2010ghosts} was somehow related to the highly redundant geometry of the WSRT, but the mechanism by which they arose was not clear. Another conclusion of that study was that ghosts could be minimized and eventually driven below the noise by laboriously improving the sky model and/or applying direction-dependent solutions.  The problem of calibration artefacts is becoming more important with the advent of new observational techniques and new radio telescopes such as the Low Frequency Array (LOFAR), the upgraded Jansky Very Large Array (JVLA), MeerKAT, etc., as well as the upcoming Square Kilometer Array (SKA). Not only do these telescopes promise (and in some cases already deliver) unprecedented sensitivity, they also increase the data rates substantially. The increased sensitivity means that fainter artefacts cannot be ignored, while the data rates require that calibration become largely automated, with careful and laborious manual data reduction no longer an option. Observations of the diffuse HI in the $0.5<z<20$  in order to probe galaxy formation \\citep[e.g.][]{EOR2006} or the nature of dark energy \\citep{IM2010} need to face exquisite calibration in order to subtract foreground that are spatially and spectrally orders of magnitude brighter \\citep[i.e.][]{Bernardi-EOR-foregrounds,Switzer-IM,Pober-PAPER-foregrounds}. Similar calibration requirements are needed for future stacking techniques for HI detection \\citep[e.g.][]{Stacking2013} and continuum surveys aimed at revealing the nJy population \\citep{Norris-SKA-continuum}. On the other hand, transient source detection pipelines require very fast on-the-fly calibration, which necessitates the use of shallow and incomplete sky models (Wijnholds 2013, priv. comm.) It is clear that a deeper and theoretical understanding of calibration artefacts, including those buried in the noise, is required if these new instruments and techniques are to achieve their scientific goals.  We have been investigating two classes of calibration artefacts, namely spurious point sources (ghosts) and source suppression (i.e. the reduction in observed flux of sources not included in the sky model). This work establishes that both are manifestation of the same underlying mechanism, and aims to provide a theoretical understanding of this. The objective is to extend the results of previous papers \\citep{Linfield1986,Wilkinson1988,Marti2008,Taylor1999}, by studying the underlying theoretical principles which are responsible for ghost formation and source suppression (as opposed to conducting an empirical survey of the different types of calibration artefacts that have been identified).  This paper is the first in a series; ongoing work is concentrating on ghost formation and source suppression in non-redundant interferometers, and on direction-dependent calibration. Ghosts have already been spotted in LOFAR data (de Bruyn priv. comm.), and patricularly in the presense of transient sources (Fender priv. comm.), and early indications are that the same mechanism is responsible. Future papers in the series will cover these phenomena. ",
        "conclusions": "In this work, we have demonstrated that the ghost source phenomenon is a fundamental aspect of the selfcal process, and will invariably arise during calibration with an incomplete sky model. It is perhaps surprising that these features were not explicitly noted before 2004, but radio interferometrists are accustomed to esoteric instrumental artefacts showing up in poorly-calibrated maps, so something that is faint to begin with and goes away with a more accurate calibration would not necessarily have attracted attention. The strongest ghost -- the one sitting on top of the missing model source -- has been indirectly observed as a matter of course, under the guise of ``flux suppression'', but being masked by the source itself, it did not give enough clues to the true extent of the phenomenon. Consider also that the spacing between ghosts (given by $\\bmath{s}_0/\\phi_0$) decreases as $\\bmath{s}_0$ becomes small, while the size of the ghosts remains the same (as it is similar to the PSF size). There is therefore some critical $\\bmath{s}_0$ beyond which the ghosts will begin to blend together into a single ``spoke'' connecting the sources. In the authors' experience, poorly calibrated WSRT maps will often exhibit spokes connecting nearby brighter sources, which tend to go away once the sky model is improved. It is quite plausible that we've been seeing ``blended ghosts'' all along. Since the primary beam makes it more likely for apparently bright sources to be near rather than distant, such blends may be the most frequent manifestation of the ghost mechanism.  Although we have only studied the two-source case theoretically, empirical results (as well as the observations of Fig.~\\ref{fig:2010ghosts}) suggest that with multiple missing  sources (which would be a far more typical case), many individual ``ghost strings'' are produced at different orientations.  This would cause the artefacts and their positive/negative sidelobes to overlap and, to an extent, average out, thus washing out the regular geometric pattern and further contributing to their non-detection. It took a special set of circumstances --  i.e. a single relatively strong and distant unmodelled source, high sensitivity, and WSRT's regular and highly redundant geometry  -- to make the ghosts stand out (Fig.~\\ref{fig:2004ghosts}) as a regular and peculiar geometric feature.  Given the increased sensitivity of current and upcoming instruments, and the emergence of statistical detection techniques, as argued in the introduction, a good understanding of the ghost phenomenon is vital. We have developed a theoretical framework for explaining ghost formation in the two-source WSRT case, which has yielded predictions that closely match simulations, qualitatively match actual observations (which are of course always more complex than just two sources), and explain all the previously puzzling features of the ghost phenomena that were empirically observed earlier (positioning at rational fractions along a line connecting two sources, a ``ghost spread function'' that differs from the PSF, independence -- to first order -- on the flux of the modelled source). We have also established that the well-known phenomenon of flux suppression is actually just one manifestation of the same underlying mechanism.  A particularly intriguing avenue of further research is ghost formation in the presense of direction-dependent effects. Recall that in the original QMC result (Fig.~\\ref{fig:2010ghosts}), the ghosts were found to be associated with uncorrected DDEs, and went away once direction-dependent solutions had been applied. In this case, the unmodelled flux responsible for ghost formation must have been due to errors in the voltage beam, caused by the artifically large pointing error. We may postulate that any unmodelled flux, whether due to errors in the sky model, or unaccounted-for DDEs, will lead to ghosts on some level.  More work is required to describe the ghost phenomenon, both empirically and theoretically, for 2D arrays, for multiple and extended unmodelled sources, and for various forms of direction-dependent calibration. In particular, flux suppression, as the strongest manifestation of the ghost mechanism, needs to be studied in more detail (especially in light of upcoming deep, blind surveys). We must also investigate alternative calibration approaches. In particular, robust calibration \\citep{Kazemi2013a} has been shown to result in less flux suppression, and must necessarily exhibit different ghost behaviour.  Ultimately, we need to develop a theoretical and/or numerical mechanism for answering the following fundamental questions. Given an observational scenario, how deep/accurate does a sky model need to be in order to suppress ghosts to a given level? And, what then are the statistical properties and signatures (power spectrum, etc.) of the remaining ghost artefacts? Building on the theory and numerical tools developed in this work, prospects are good that we can eventually provide rigorous answers to these questions."
    },
    "1402/1402.4760_arXiv.txt": {
        "abstract": "We propose a possible new mechanism for a strong and stable magnetic field of compact stars due to an instability in the presence of a chirality imbalance of electrons---the chiral plasma instability. A large chirality imbalance of electrons inevitably occurs associated with the parity-violating weak process during core collapse of supernovae. We estimate the maximal magnetic field due to this instability to be of order $10^{18}$ G at the core. This mechanism naturally generates a large magnetic helicity from the chiral asymmetry, which ensures the stability of the large magnetic field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.2370_arXiv.txt": {
        "abstract": "We have investigated the influence of velocity shear and a radial density profile on the spatial development of the current driven kink instability along helically magnetized relativistic jets via three-dimensional relativistic magnetohydrodynamic simulations. In this study, we use a non-periodic computational box, the jet flow is initially established across the computational grid, and a precessional perturbation at the inlet triggers growth of the kink instability.  If the velocity shear radius is located inside the characteristic radius of the helical magnetic field, a static non-propagating current driven kink is excited as the perturbation propagates down the jet. Temporal growth disrupts the initial flow across the computational grid not too far from the inlet. On the other hand, if the velocity shear radius is outside the characteristic radius of the helical magnetic field, the kink is advected with the flow and grows spatially down the jet. In this case flow is maintained to much larger distances from the inlet. The effect of different radial density profiles is more subtle. When the density increases with radius, the kink appears to saturate by the end of the simulation without apparent disruption of the helical twist. This behavior suggests that relativistic jets consisting of a tenuous spine surrounded by a denser medium with a velocity shear radius outside the radius of maximum toroidal magnetic field have a relatively stable configuration. ",
        "introduction": "Relativistic jets are ubiquitous features of many accreting black hole and neutron star systems. They are associated with X-ray binaries (XRBs) and Gamma-Ray Bursts (GRBs) containing neutron stars or stellar-mass black holes (BHs) of mass $M \\sim 1.4 - 20~M_{\\odot}$, and active galactic nuclei (AGNs) with super massive BHs of mass $M \\sim 10^6 - 10^{10}~M_{\\odot}$.  It is commonly believed that jets are powered and collimated by magnetohydrodynamic (MHD) processes (Lovelace 1976; Blandford 1976, 2000; Blandford \\& Znajek 1977). General relativistic magnetohydrodynamic (GRMHD) simulations of jet formation (e.g., De Villiers et al. 2003, 2005; Hawley \\& Krolik 2006, McKinney \\& Gammie 2004; McKinney 2006; McKinney \\& Blandford 2009; Beckwith et al. 2008; Hardee et al. 2007; Komissarov \\& Barkov 2009; Penna et al. 2010) show development of turbulence via magneto-rotational instability (MRI) (Balbus \\& Hawley 1998) and angular momentum transfer in the accretion disk, leading to diffusion of matter and magnetic field inwards, and generation of unsteady outflows. In general, GRMHD simulations with spinning black holes indicate jet production consisting of a Poynting dominated (in the sense that the energy is transferred predominantly by the electromagnetic field) high Lorentz factor spine, and a matter dominated, mildly relativistic sheath with $v < c$ possibly embedded in a sub-relativistic speed, $v \\ll c$, disk/coronal wind.  Note however, the that the Lorentz factor can achieve a maximum in a tenuous boundary layer between a Poynting-flux spine and a dense lower speed kinetically dominated sheath via rarefaction acceleration (e.g, Aloy \\& Rezzolla 2006; Mizuno et al. 2008; Zenitani et al. 2010a; Komissarov et al. 2010).  It is thought that relativistic jets on the largest scales are kinetically dominated, i.e., most of the electromagnetic energy is converted to kinetic energy. The means by which magnetic energy is converted to kinetic energy has not been identified, but it is generally thought to involve gradual acceleration by magnetic forces (Beskin \\& Nokhrina 2006; Komissarov et al 2007, 2009, 2010; Tchekhovskoy et al. 2009, 2010; Lyubarsky 2009, 2010a, 2011; Granot et al. 2011; Lyutikov 2011; Granot 2012a,b) and/or non-ideal MHD effects, specifically magnetic reconnection (Spruit et al. 2001; Drenkhan 2002, Drenkhahn \\& Spruit 2002;  Lyubarsky 2010b; McKinney \\& Uzdensky 2012). The large scale magnetic field may dissipate if the regular magnetic structure is destroyed as a result of a global MHD instability, the kink instability being the most plausible candidate (Lyubarskii 1992, 1999; Eichler 1993; Spruit et al 1997; Begelman 1998; Giannios \\& Spruit 2006).  The most promising MHD models for the acceleration and collimation of jets involve the presence of a magnetic field with foot points anchored to a rotating object (an accretion disk or a spinning neutron star or black hole). The dominance of the toroidal component ($B_\\phi$) over the poloidal component ($B_p$) is a natural consequence in these models. It is well known that current carrying plasma columns containing strong toroidal magnetic fields are unstable to non-axisymmetric perturbations (Bateman 1978). Among these current-driven ({\\bf CD}) instabilities, the kink mode is the most violent. The kink mode leads to helical displacement of the plasma-column and may disrupt the system. Thus, non-linear development of the kink mode could trigger violent magnetic dissipation in relativistic outflows. McKinney \\& Blandford (2009) pointed out that  when the Kruskal-Shafranov (KS) instability criterion, $|B_{p}/B_{\\phi}| < z / 2 \\pi R$, for cylindrical force-free equilibria is applied to relativistic jets, the jets become unstable beyond the Alfv\\'en surface located at $z \\simeq 10 R_g$ where $R_g$ is the gravitational radius, and where the jet radius $R_j \\lesssim z$ and  $|B_\\phi| \\gtrsim |B_p|$. However, jet rotation and velocity shear significantly affect the instability criterion (Istomin \\& Pariev 1994, 1996;  Lyubarskii 1999; Tomimatsu et al. 2001; Narayan et al. 2009) and need to be studied in this context.  On the smallest scales, the rapid observed variability of X-ray/TeV gamma-ray flares with timescales from several minutes to a few hours in blazars pose severe constraints on the particle acceleration timescale and the emission region size and also provide indirect evidence for instabilities. Here particle acceleration provided by current driven instability and magnetic reconnection conversion of Poynting flux to kinetic flux (Sikora et al. 2005) to generate flares would have to come from small, a few Schwarzschild radii in size, fast moving ``needles\" or ``jet\" within a slower jet medium, the so-called ``needles or jet-in-a-jet\" scenarios (Levinson 2007; Begelman et al. 2008; Ghisellini \\& Tavecchio 2008; Giannios et al. 2009).  On somewhat larger scales, high resolution VLBI observations show that many AGN jets display helical structures on sub parsec to kiloparsec scales (e.g., G\\'{o}mez et al. 2001; Lobanov \\& Zensus 2001) that provide indirect evidence for instability. Helical distortion may be caused by precession of the jet and/or by MHD instabilities (current-driven or Kelvin-Helmholtz velocity shear driven). Non-relativistic and relativistic simulations of magnetized jet formation and/or propagation have showed helical structures attributed to CD kink instability (e.g., Lery et al. 2000; Ouyed et al. 2003; Nakamura \\& Meier 2004; Nakamura et al. 2007; Moll et al. 2008; Moll 2009; McKinney \\& Blandford 2009; Carey \\& Sovinec 2009; Mignone et al. 2010; Porth 2013; Guan et al. 2013). In the absence of CD kink instability and resistive relaxation, helical structures may be attributed to the Kelvin-Helmholtz ({\\bf KH}) helical instability driven by velocity shear at the boundary between the jet and the surrounding medium (e.g., Hardee 2004, 2007; Mizuno et al. 2007) or triggered by precession of the jet ejection axis (Begelman et al. 1980). It is still not clear whether CD, velocity shear driven, or jet precession is responsible for the observed structures, or whether these different processes are responsible for the observed twisted structures at different spatial scales.  In previous work, we investigated the temporal development of the CD kink instability using a periodic computational box. In Mizuno et al. (2009), we studied the instability of helically magnetized static plasma columns (or more generally rigidly moving flows considered in the proper reference frame). Simulation results showed that the initial configuration was strongly distorted but not disrupted by the CD kink instability. The growth rate depended on the radial profile of the magnetic pitch and the density. The plasma column was more unstable when the helicity of the magnetic field lines increased with radius (magnetic pitch parameter, $P \\equiv RB_z/B_{\\phi}$, decreased with radius $R$) and when the density decreased with radius. In Mizuno et al. (2011), we investigated the influence of velocity shear on the CD kink instability and found that the kink propagates along the jet with speed and flow structure dependent on the radius of the velocity shear (hereafter called the velocity shear radius) relative to the characteristic radius of the helically twisted force-free magnetic field. In Mizuno et al. (2012), we studied the influence of jet rotation and differential motion on the CD kink instability, and in accordance with linear stability theory, we found that development of the instability depends on the lateral distribution of the poloidal magnetic field. When the profile of the poloidal field is shallow, the instability develops slowly and eventually saturates. When the profile of the poloidal field was steep and the magnetic pitch in the inner portion of the jet was sufficiently small, i.e., the magnetic field was sufficiently tightly twisted, multiple growing helical wavelengths disrupted the initial configuration.  Recently O'Neill et al. (2012) have also investigated the CD kink instability in a local, co-moving frame using the 3D RMHD module within the Athena code (Beckwith \\& Stone 2011). They also studied rotating jets with a purely toroidal magnetic field, the hoop stress being balanced by the gas pressure and centrifugal force, and found that the instability brings the system into a turbulent state. All of these results have been obtained using a periodic computational box that allows study of the temporal growth of a few axial wavelengths that fit within the small simulation domain.  The time has come to begin systematic more physically realistic simulations that study the spatial development of CD instability.  In order to investigate the spatial development of the CD kink instability a non-periodic computational box must be used. At this time, only a few 3D simulations of relativistic, magnetized jets have been performed that study the spatial properties of the CD kink instability, and the results are still controversial. McKinney \\& Blandford (2009) simulated the generation and propagation of a relativistic highly magnetized jet and found no significant development of non-axisymmetric instabilities, even though the simulation satisfied the KS criterion for instability. On the contrary, Mignone et al. (2010) found strong jet distortion in their jet simulation. The discrepancy may be attributed to the difference in the setup. Mignone et al. (2010) assumed that the magnetic field was purely toroidal, whereas McKinney \\& Blandford (2009) simulated jet launching from a rotating magnetized configuration and their jet contained both toroidal and poloidal magnetic field components. Recently, Porth (2013) performed a simulation of relativistic jet formation from magnetospheres in Keplerian rotation. The simulation showed development of the kink mode but the kink saturated before notable dissipation or disruption occurred.  In this paper, we begin a systematic study of the spatial development of the CD kink instability in a relativistic jet via 3D RMHD simulations using a non-periodic computational box. In order to isolate the role of different factors and therefore to gain physical insight we do not attempt to properly account for poorly known relativistic jet poloidal and radial profiles. Here we concentrate on the role of the velocity shear radius relative to the radius at which the toroidal component of the magnetic field is maximum and the radial density profile. We do not investigate the effect of a radial velocity or magnetic pitch profile in this study. We know that the temporal development of the CD kink instability is dependent on differential motion, the magnetic pitch profile, and the density profile (Appl et al. 2000; Mizuno et al. 2009, 2011, 2012).  Our previous temporal growth simulations (Mizuno et al. 2009) considered declining and constant density profiles with constant density profiles being slightly more stable. In order to compare spatial results to previous temporal results and to more closely simulate a high speed tenuous jet surrounded by a denser sheath or wind, we consider declining, constant and increasing density profiles. Since temporal stability results (Mizuno et al. 2009, 2011, 2012) indicate more substantial instability dependence on the magnetic pitch and velocity profiles that are unknown for observed relativistic jets, a choice of constant magnetic pitch and sharp velocity shear is reasonable for this first study.  We describe the numerical method and setup used for our simulations in Section 2, present our results in Section 3, and in Section 4 compare our results to instability expectations and conclude. ",
        "conclusions": "We have investigated the influence of velocity shear location and radial density profile on the spatial development of the CD kink instability using a non-periodic computational box. In these spatial growth simulations and the previous temporal growth simulations (Mizuno et al.\\ 2011) the mildly relativistic jet is established across the computational domain with a helical force-free magnetic field. Our previous study of temporal growth of the CD kink instability used a periodic computational box and uniform spatial perturbation of the jet across the box that only allowed for temporal development of wavelengths that fit within the box. The kink propagated or was static depending on the velocity shear radius relative to the characteristic radius of the magnetic field. For the case of constant helicity and constant magnetic pitch that we have considered here, the characteristic radius corresponds to the radius at which the toroidal magnetic field is a maximum.  For temporal growth cases studied previously using a periodic computational box (Mizuno et al. 2011) with increasing or decreasing helicity, i.e., for decreasing or increasing magnetic pitch, the characteristic radius is outside or inside the radius at which the toroidal magnetic field is a maximum, respectively, but not by a large amount.  Our temporal studies (Mizuno et al.\\ 2011) showed that results are strongly affected by the choice of velocity shear radius with respect to the characteristic radius of the helical magnetic field. Simulations found a temporally growing static kink for a velocity shear radius of $R_j = a/2$ that is the same as the velocity shear radius of our present cases (A1, B1, C1). Kink propagation was found to increase as the velocity shear radius increased and found a temporally growing kink propagating at about the jet speed for a velocity shear radius of $R_j = 4a$ that is the same as the velocity shear radius of our present cases (A2, B2, C2). Other temporal studies investigated the influence of jet rotation and differential motion on the CD kink instability (Mizuno et al.\\ 2012). The most unstable of these temporal simulations developed both longer and shorter kink wavelengths at larger and smaller jet radii, respectively, and this combination led to disruption of the initial cylindrical configuration.  The shorter wavelength inner kink developed when the wavelength at one-half the periodic computational box length became unstable as a result of increasing the magnetic helicity (decreasing the magnetic pitch) in the inner portion of the jet. While this decrease in the magnetic pitch was associated with jet rotation, in general the temporal growth properties were not much different than for the comparable non-rotating static cases studied in Mizuno et al. (2009).  In the present simulations, we find that the CD kink instability is partially stabilized by a radially increasing density structure. Current observations and GRMHD simulation of jet formation suggest that an inner high speed jet is surrounded by a denser and slower outer flow (e.g., McKinney 2006). Therefore, we expect that in the GRMHD simulation by McKinney \\& Blandford (2009), a CD kink develops in the jet but saturates in part as a result of radially increasing density structure. This stabilizing effect helps us to understand the existence of non-destructive kink structure in observed relativistic jets. The kink (wiggle) structure is developed by the CD instability near to the jet origin but is sufficiently stabilized at larger distances by the radially increasing density structure and other effects such as jet expansion, axial velocity and rotation profiles, etc. and we will investigate these effects in future work.  Our simulations indicate that CD kink wavelength(s) and propagation dependent on the location of velocity shear, magnetic pitch profile and characteristic magnetic radius. A characteristic magnetic radius inside the velocity shear combined with a less dense jet inside a denser confining medium is stabilizing. In order to make any comparison with observed helical structures we need to perform radiation transfer calculations based on simulation results making image and polarization maps to directly compare with observations. Ultimately the combination of simulation generated images and correspondingly detailed observation images might allow the determination of AGN jet radial velocity and magnetic field profiles. Existing VLBI observations indicate moving helical jet structures suggesting an effective velocity shear radius on the order of or greater than the characteristic radius of the magnetic field. To further resolve this issue would require VLBI resolution sufficient to make a comparison between polarization observations giving radial magnetic structure and proper motion observations giving radial velocity structure. Such futuristic observations could be used to pin down where the velocity is most strongly sheared compared to the characteristic radius of the jet's magnetic field.  In this study, we have used an initially constant axial speed with a non-rotating but CD kink unstable setup as a first step in our investigation of spatial growth of the CD kink instability. The precessional perturbation that we impose at the inlet winds up the initially non-rotating force-free magnetic field configuration leading to jet acceleration and rotation. Interestingly, we find a reversal in the sense of rotation across an axial velocity minimum that is well inside the velocity shear radius. Such a reversal in rotation is not seen in the propagating jet simulations performed by Porth (2013) although his simulations have shown faster rotation accompanying the location of stronger toroidal magnetic fields wound up by the inlet rotation profile. These profile differences show, as found by O'Neill et al.\\ (2012), that structural details are significantly dependent on the initial configuration. A more realistic setup for relativistic jets would involve more physically relevant radial toroidal and poloidal velocity and magnetic profiles along with jet acceleration and expansion. Jet acceleration and expansion change the behavior of the instabilities seen on jets modeled as cylinders of constant radius. Jet expansion has a spatially stabilizing effect as spatial growth scales with the characteristic magnetic and jet radius. Additionally, magnetic instabilities become ineffective when the Alfv\\'{e}n speed drops below the lateral expansion speed of the jet. Currently only a few MHD simulations have explored the effect of jet expansion on the development of the CD kink instability (Moll et al. 2008; Moll 2009; Porth 2013). In future work, we plan to investigate the effect of jet acceleration and expansion on the development of CD kink instability in the relativistic MHD regime along with more realistic radial profiles.  We have used a mildly relativistic jet speed, $v_j=0.2c$, in the simulations. This has the advantage of allowing  comparison to non-relativistic 3D MHD simulations and will provide a baseline for comparison with relativistic jet speed simulations. In general, our present results are similar to those obtained from non-relativistic 3D MHD simulations of jet formation and propagation (e.g., Nakamura \\& Meier 2004; Moll et al. 2008; Moll 2009). For example, exponential growth times in our simulations and previous non-relativistic MHD simulations are generally on the order of Alfv\\'{e}n crossing times. In non-relativistic MHD simulations of jet propagation, the helically twisted structures developed by CD kink instability move with the jet propagation speed (e.g., Nakamura \\& Meier 2004; Moll et al. 2008; Moll 2009), and we find similar results provided the jet's velocity shear radius is outside the characteristic radius of the helical magnetic field. The advantage to our setup, unlike jet generation and propagation based simulations, is control of radial and axial profiles including acceleration and expansion. This will allow a systematic determination of the most important stabilizing factors.  In our current and past simulation work, we assume ideal MHD with no resistivity other than induced by the numerical scheme. Thus, relativistic magnetic reconnection (e.g., Blackman \\& Field 1994; Lyubarsky 2005; Watanabe \\& Yokoyama 2006; Zenitani et al. 2010b; Takahashi et al. 2011; Zanotti \\& Dumbser 2011; Baty et al. 2013; Takamoto 2013) associated with non-linear kink development is not treated. Development of resistive relativistic MHD (RRMHD) codes are now one of the frontiers of relativistic MHD (e.g., Komissarov 2007; Palenzuela et al. 2009; Dumbser \\& Zanotti 2009; Zenitani et al. 2010b; Takahashi et al. 2011; Takamoto \\& Inoue 2011). We have developed an RRMHD code (Mizuno 2013). The present and ongoing work using ideal MHD will be used to establish a baseline for comparison with CD kink instability work allowing for relativistic magnetic reconnection."
    },
    "1402/1402.2236_arXiv.txt": {
        "abstract": "We test the FLRW cosmology by reconstructing in a model-independent way both the Hubble parameter $H(z)$ and the comoving distance $D(z)$ via the most recent Hubble and Supernovae Ia data. In particular we use: data binning with direct error propagation, the principal component analysis, the genetic algorithms and the Pad\\'e approximation. Using our  reconstructions we evaluate the Clarkson {\\it et al} test known as $\\Omega_K(z)$, whose value is constant in redshift for the standard cosmological model, but deviates elsewise. We find good agreement with the expected values of the standard cosmological model within the experimental errors. Finally, we provide forecasts, exploiting the Baryon Acoustic Oscillations measurements from the Euclid survey. ",
        "introduction": "The observed accelerated expansion of the late-time universe, as evidenced by a number of cosmological data like type Ia supernovae (SnIa) \\cite{sn1a}, the cosmic microwave background radiation \\cite{cmb} (CMB) and large scale structure \\cite{lss} (LSS) came as a great surprise to cosmologists. It is quite straightforward to explain the effect within the framework of Friedmann-Lema\\^itre-Robertson-Walker (FLRW) cosmology by simply introducing a cosmological constant or a more general (dynamical) dark energy component. However, all such components give rise to severe coincidence and fine-tuning problems. An alternative approach postulates that General Relativity is only accurate on small scales and modifications at larger scales are needed. These modification would lead to the observed late-time acceleration \\cite{CD,BDEL,RM,fR}.  Recently, several works have tried to take a different approach, namely to study the effect of large scale structure on the observed luminosity-to-redshift behavior of SnIa (the first observable that has led to the conclusion that the Universe is accelerating). The assumption that the Universe is homogeneous and isotropic is known to be violated at late times. However, the scale at which the assumption breaks down is still in debate, \\cite{campusano}. Nonetheless, recent studies have been able to reproduce the luminosity-to-redshift relation provided that we live in a large region, called empty void, in which the matter density is less than the spatial average density on large cosmological scales (\\cite{LTB}, for a recent review see \\cite{voids-review}). However, also these models are still in debate since they require the observer to be located at the center of the void or at most within about few percent of the void scale radius \\cite{BM}, thus disfavoring the model from a Copernican Principle (CP) point of view.  Despite all the theoretical efforts to understand the accelerated expansion of the Universe, see \\cite{ks, sapone}, right now there is no model capable of providing singlehandedly a satisfactory explanation. An alternative is to create null tests for the current paradigm and see if and how it breaks down. For example, Clarkson {\\it et al.} \\cite{cbt} presented a new test for the CP which relies on a consistency relation that exists within the homogenous and isotropic FLRW model between the luminosity and the Hubble parameter. This relation is expected to hold exactly at all redshifts $z$, and any deviation will point to a departure from the FLRW model. Indeed, to be precise, the test of Clarkson {\\it et al.} tests deviations from a FLRW metric and not strictly from the CP, as pointed out by \\cite{wil2, wil, s1}.  The existence of this consistency relation implies that if we have two separate experiments measuring independently the comoving distance $D(z)$ (or equivalently the angular diameter distance or the luminosity distance) and the Hubble parameter $H(z)$, then we are able to reconstruct the curvature parameter $\\Omega_{K}$ at each redshift in a model-independent way. In a FLRW universe, the curvature parameter should not depend on redshift, hence $\\Omega_{K}$ should not vary, so if we measure a variation of the curvature parameter over redshift, this means that the assumption of homogeneity at large scales has to be rejected.  Since it was proposed, this test has raised interest in the astrophysics community. Ref. \\cite{avg, sha} have applied it by using available data; ref. \\cite{s4} has discussed the deviation from the FLRW relation due to backreaction; refs. \\cite{wil, feb, s2, s3} have predicted the outcome of the test, i.e. the behavior of $\\Omega_K$ with redshift, in the case of their particular inhomogeneous cosmology or void model; ref. \\cite{lar} has made predictions on future results from the test if a toy-model of backreaction is assumed to be the correct cosmology; ref. \\cite{mor} has used it to measure the curvature of the Universe in a model-independent way. Also, many other tests of the standard model of cosmology have been outlined (for most recent work, see \\cite{s5}, and for a review see \\cite{Clarkson-review, Marra-review})  In particular, ref. \\cite{avg} has used $H(z)$ data from passively evolving galaxies \\cite{Jimenez:2003iv, Simon:2004tf} and the SnIa of \\cite{Kowalski:2008ez} and has found no indication of a deviation from FLRW. Ref. \\cite{mor} has used SnIa data of \\cite{Amanullah:2010vv}, H(z) data from passively evolving galaxies \\cite{Stern:2009ep} and baryonic acoustic oscillation (BAO) data \\cite{Gaztanaga:2008xz}, and the $H_0$ measurement of \\cite{Riess:2009pu}. Similarly, ref. \\cite{sha} has used the SnIa of \\cite{Hicken:2009dk} and the measurement of $H_0$ of \\cite{Riess:2009pu} or of \\cite{Komatsu:2008hk}, in combination with $H(z)$ measurements coming either from passively evolving galaxies \\cite{Jimenez:2003iv, Stern:2009ep}, or from BAO data \\cite{Percival:2009xn}, together with cosmic microwave background data \\cite{Wang:2007mza}. No evidence for deviations from FLRW nor from flatness was found.  In this context, our work has two objectives, building on the work of \\cite{sha} and \\cite{lar}. The first objective is to improve the measurement of $\\Omega_K$ by comparing four different measurement techniques based both on binnings of the data and on very efficient model-independent reconstructions of functions, and by using the most recent data. The second is to make new up-to-date forecasts for the $\\Omega_K$ test for future LSS and SnIa data.  The paper is organized as follows: in Sec.~\\ref{sec:background} we describe the general background equations used for our analysis. In Sec.~\\ref{sec:reconstruction_H_D} we describe four different methods to reconstruct the Hubble parameter and the comoving distance from two different data sets. In particular we show: the data binning and direct error propagation method, the principal component analysis (PCA), the genetical algorithms (GA) and the Pad\\'e approximation. In each subsection we report the results. Sec.~\\ref{sec:forecasts} is devoted to the Fisher matrix forecasts on the errors for the $\\Omega_{K}$ at different $z$ using the Euclid galaxy redshift survey\\footnote{http://www.euclid-ec.org/} \\cite{Cimatti:2009is, RedBook} and future SnIa surveys. ",
        "conclusions": "}  In this paper we used four different methods to reconstruct, in a model-independent fashion, the Hubble parameter and the luminosity distance, measured by two different data sets. We find that all four methods used, i.e. direct binning, principal component, genetic algorithms and Pad\\'e approximation, give results that are in agreement with respect to their general behavior.  In particular, all four methods lead to very large errors at low redshifts, while errors decrease noticeably for larger $z$. The transition redshift is in all cases between $0.2$ and $0.4$. To be specific, as for the direct binning and the PCA techniques, the best measured $\\Omega_K$ have errors $\\sigma_{\\Omega_K} \\simeq 3$. For the genetic algorithms and Pad\\'e approximations instead, we found that the errors on the curvature parameter are  of the order of about $\\sigma_{\\Omega_K}\\sim 20$ at small redshifts, and they decrease when $z$ grows, reaching a value of about $\\sigma_{\\Omega_K}\\sim 0.5$ for redshifts $z\\geq1$. This means that reconstruction techniques manage to reduce the error if compared to techniques where data need to be binned. The best method, leading to the smallest error on $\\Omega_K$, $\\sigma_{\\Omega_K}\\simeq 0.1$, is the one based on GA.  We also find that the choice of the binning, as expected, influences the values of the errors, by affecting mostly the derivative of the comoving distance. This is simply due to the nature of the SnIa data: in regions where the values of the distance modulus have large fluctuations, there will be a strong dependence of the choice of binning on the value of the derivative.  Going back to the problem of the large size of the errors at low redshifts, this may appear strange because here we have a large number of data with small errors (in particular as regards SnIa data, see e.g. Fig. \\ref{fig:comoving-distances}). If we have a closer look though, we find that the main reason for this feature lays on the expression of $\\Omega_K$ itself, so let us explain this in detail. The curvature parameter was found to be \\be \\Omega_{K} = \\frac{\\left[H(z)D_{,z}(z)\\right]^2-1}{\\left[H_0D(z)\\right]^2}\\,. \\label{eq:omegak-concl} \\ee This expression tells us that if the Universe is not homogeneous then the curvature parameter has to vary with redshift, i.e. $\\Omega_K$ is not constant. From Eq.~(\\ref{eq:curvAngDist}) we learn that the comoving distance is the integral of the Hubble parameter; however, Eq.~(\\ref{eq:curvAngDist}) has been evaluated under the assumption of homogeneity, as a consequence $\\Omega_K$ can only be constant.  Now let us assume that, via one observable, we measure the Hubble parameter $H(z)$ and, via another independent observable, we measure the luminosity distance, which gives a slightly different Hubble parameter, say $H_1(z)$. Then, substituting $H_1(z)$ in Eq.~(\\ref{eq:curvAngDist}), Eq.~(\\ref{eq:omegak-concl}) will look like: \\be \\tilde{\\Omega}_{K}(z) = \\Omega_{K}\\frac{1 -\\left(\\frac{H(z)}{H_1(z)}\\right)^2\\cos^2\\left[\\sqrt{-\\Omega_K}\\int_{0}^{z}{\\frac{H_0{\\rm d}x}{H_1(x)}}\\right]}{\\sin^2\\left[\\sqrt{-\\Omega_K}\\int_{0}^{z}{\\frac{H_0{\\rm d}x}{H_1(x)}}\\right]}\\,, \\label{eq:omegak-lowz} \\ee where $\\tilde{\\Omega}_K$ is the function that we want to evaluate, whereas $\\Omega_K$ is the standard curvature parameter. As we can see, for small redshifts, the numerator of Eq.~(\\ref{eq:omegak-lowz}) stays finite ($\\cos(0) =1$), but the denominator does not ($\\sin(0)=0$). Let us assume that for $z\\rightarrow 0$, the quantity $\\sqrt{\\Omega_K}\\int_{0}^{z}{H_0{\\rm d}x/H_1(x)}\\rightarrow \\epsilon$, where $\\epsilon$ is a small number that goes to zero when $z = 0$. Then, Eq.~(\\ref{eq:omegak-lowz}) reads \\be \\tilde{\\Omega}_{K}(z) = \\Omega_{K}\\frac{1-\\left(\\frac{H(z)}{H_1(z)}\\right)^2\\cos^2\\left(\\epsilon\\right)}{\\sin^2\\left(\\epsilon\\right)}= \\Omega_{K}\\frac{1-\\left(\\frac{H(z)}{H_1(z)}\\right)^2+\\left(\\frac{H(z)}{H_1(z)}\\right)^2\\sin^2\\left(\\epsilon\\right)}{\\sin^2\\left(\\epsilon\\right)} \\label{eq:omegak-lowz1} \\ee For $\\epsilon \\rightarrow 0$, Eq.~(\\ref{eq:omegak-lowz1}) becomes \\be \\tilde{\\Omega}_{K}(z)=\\lim_{\\epsilon \\rightarrow 0}\\left(\\Omega_{K}\\frac{1-\\left(\\frac{H(z)}{H_1(z)}\\right)^2+\\left(\\frac{H(z)}{H_1(z)}\\right)^2\\sin^2\\left(\\epsilon\\right)}{\\sin^2\\left(\\epsilon\\right)} \\right) \\rightarrow \\infty\\,, \\ee Then, in order for the above equation to stay finite at small redshift we need $H_1(z)$ to be exactly equal to $H(z)$; in other words, the curvature parameter stays constant at small redshifts if the Universe is homogeneous! This is the reason why for small redshifts, when measuring $H(z)$ and $D(z)$ with independent observations, the values of the curvature parameter (and their errors) are very large. In this paper, we used two different datasets which give, of course, two different best fits of the parameters and consequently these two best fits will give two different Hubble parameters ($H(z)$ and $H_1(z)$), so that Eq.~(\\ref{eq:omegak-lowz1}) still diverges when $z\\rightarrow 0$.  Alternatively, we can also perform a series expansion on Eq.~(\\ref{eq:omegak-concl}) for small $z$. Doing so we find \\be \\tilde{\\Omega}_{K}=\\Omega_k+\\frac{1-\\frac{H_{1,0}^2}{H_0^2}}{z^2}+\\frac{\\frac{2 H_0 H'(0)-H_{1,0} H_1'(0)}{H_0^2}-\\frac{H_1'(0)}{H_{1,0}}}{z}+..., \\ee where the primes denote a derivative with respect to $z$, eg $H_1'(0)=\\frac{dH_1}{dz}|_{z=0}$, while $H_0$ and $H_{1,0}$ are the values of the two Hubble parameters at $z=0$. The advantage of this approach is that we can also clearly see, in a model-independent way and without assuming a Dark Energy model, the type of the singularity as $z\\rightarrow0$. So, unless $H(z)=H_1(z)$ then there will be a singularity $\\sim1/z^2$ at $z=0$.  To also prove numerically what is explained above, we report the case of the PCA analysis. Here the value of the Hubble parameter  in each of the 6 bins, obtained when using the best fits from SnIa are: $H(z_i) =\\{71.1431, 76.9306, 87.9404, 103.691, 108.447, 129.132\\}$. When using the best fits from $H(z)$ measurements from passively evolving galaxies, we obtain instead: $H(z_i) =\\{71.1061, 82.0685, 91.8856, 102.916, 131.697, 166.138\\}$. The difference of the Hubble parameters in the first redshift bin is $0.037$; even though this number is very small, it is not small enough to guarantee the curvature parameter to be stable at low redshifts.  The same conclusion was previously found by \\cite{s3}\\footnote{However, here a divergence $\\propto z^{-1}$ was found}. Also \\cite{sha} explain this feature of the $\\Omega_K$ test, but their analysis was done assuming a particular dark energy model and the results could be different for different models. Here we have demonstrated that  Eq.~(\\ref{eq:omegak-concl}) always diverges at small redshifts independently of the model as $\\sim1/z^2$, for all models different from the homogeneous and isotropic one.  As a last remark, it is important to notice that in general the data are not free of systematics (and they are unavoidable), so even if the universe is homogeneous and isotropic, we will always reconstruct two different Hubble parameters from two different datasets, which will lead to a divergent (or very large) $\\Omega_K$ at low redshifts. Probably the test of $\\Omega_K$ will give a results that it is still consistent with the FLRW universe within the errors. Also, one should consider that there will always be a ``noise'' in the measurement of $\\Omega_K$ due to the presence of inhomogeneities, that generate fluctuations on the measured values of the pure FLRW cosmological parameters \\cite{wessel}.   Future data from the Euclid satellite will improve considerably the errors on the curvature parameter with respect to present data; in particular we will have an improvement of about $10$ times when using Euclid only, and even up to $40$ times if we add future SnIa data to the Euclid survey data."
    },
    "1402/1402.5416_arXiv.txt": {
        "abstract": "We report the discovery of WTS-2 b, an unusually close-in 1.02-day hot Jupiter ($M_{P}=1.12\\mjup$, $R_{P}=1.363\\rjup$) orbiting a K2V star, which has a possible gravitationally-bound M-dwarf companion at $0.6$ arcsec separation contributing $\\sim20$ percent of the total flux in the observed $J$-band light curve.  The planet is only 1.5 times the separation from its host star at which it would be destroyed by Roche lobe overflow, and has a predicted remaining lifetime of just $\\sim40$ Myr, assuming a tidal dissipation quality factor of $\\qprime=10^{6}$. $\\qprime$ is a key factor in determining how frictional processes within a host star affect the orbital evolution of its companion giant planets, but it is currently poorly constrained by observations. We calculate that the orbital decay of WTS-2 b would correspond to a shift in its transit arrival time of $T_{\\rm shift}\\sim17$ seconds after $15$ years assuming $\\qprime=10^{6}$. A shift less than this would place a direct observational constraint on the lower limit of $\\qprime$ in this system. We also report a correction to the previously published expected $T_{\\rm shift}$ for WASP-18 b, finding that $T_{\\rm shift}=356$ seconds after $10$ years for $\\qprime=10^{6}$, which is much larger than the estimated $28$ seconds quoted in WASP-18 b discovery paper. We attempted to constrain $\\qprime$ via a study of the entire population of known transiting hot Jupiters, but our results were inconclusive, requiring a more detailed treatment of transit survey sensitivities at long periods. We conclude that the most informative and straight-forward constraints on $\\qprime$ will be obtained by direct observational measurements of the shift in transit arrival times in individual hot Jupiter systems. We show that this is achievable across the mass spectrum of exoplanet host stars within a decade, and will directly probe the effects of stellar interior structure on tidal dissipation. ",
        "introduction": "The orbital period distribution of gas giants is a fundamental property of planetary systems and places constraints on their formation processes, migration mechanisms, and future evolution. The observed period distribution is not smooth. The majority of hot Jupiters, i.e. those with a semi-major axis $a<0.1$ AU, are found in a `pile-up' at periods of $P\\sim3-4$ days ($\\sim0.035-0.045$ AU), whereas only four hot Jupiters are found in very close in orbits ($\\lesssim0.02$ AU, $\\lesssim1$ days), namely WASP-18 b, WASP-19 b, WASP-43 b, and WASP-103 b (\\citealt{Hel09,Heb10,Hel11,Gil14}, respectively). The sharp decline of hot Jupiters in orbital periods less than two days is a genuine feature of the exoplanet period distribution, confirmed by both ground-based and space-based planet searches, e.g.  \\textit{Kepler} \\citep{How11}, and SuperWASP \\citep{Hel12}.  This suggests that very close-in hot Jupiters are relatively rare, else current instrumentation would easily detect them on account of their very frequent and deep transits, and large RV variations in comparison to longer period, smaller planets.  As a result, \\citet{Hel11} argue that extreme systems like WASP-19 b are approximately one hundred times less common than those hot Jupiters in the pile-up, indicating that it is either difficult to get gas giants into very close orbits, or that they are quickly destroyed by strong tidal forces once they arrive. The latter would imply that very close-in hot Jupiters with old host stars are in the last few percent of their lifetimes, which raises a further question of how likely it is to have observed these systems in a transient phase of their orbital evolution. Despite extensive theoretical work, our understanding of how tidal forces influence the orbital evolution of giant planets is poorly constrained by observation. The efficiency of the dissipation of the orbital energy due to frictional processes in the star is usually parameterised by a stellar tidal quality factor $\\qprime$. Studies of binary star systems estimate its value to be $\\qprime\\sim10^{6}$ (see e.g. \\citealt{Mei05}) and analysis of the tidal evolution of a small sample of exoplanets has found some evidence for consistency with this value ($10^{6}<\\qprime<10^{9}$, \\citealt{Jack08}). On the other hand, a recent exoplanet population study which tuned the value of $\\qprime$ until the distribution of remaining planet lifetimes was statistically likely, found $\\qprime\\gtrsim10^{7}$ at the $99\\%$ confidence level \\citep{Pen12} for its specific set of initial conditions. However, direct observational measurements of $\\qprime$ in individual systems, i.e. the observation of a decaying orbital period, do not currently exist.  $\\qprime$ is the dominant factor in setting the pace of the orbital evolution for very close in hot Jupiters and the unusually short predicted remaining lifetimes for planets such as WASP-18 b and WASP-19 b have lead to a number of suggested modifications to the theory of stellar tides that reduce the efficiency of the dissipation. For example, \\citet{Win10} speculate that the observed increase in misalignment between the planetary orbit and stellar spin axes for hot Jupiters orbiting hotter stars depends on the depth of the convective zone in the host star.  Here, cooler stars with deeper convective envelopes dissipate the orbital energy more efficiently resulting in a faster alignment of the stellar obliquity, in keeping with theoretical studies (e.g.  \\citealt{Bark09,Bark10,Pen11}). Others suggest that there is a complicated dependency on the planetary mean motion that results in zones of inefficient tidal dissipation during inspiral \\citep{Ogi07}, or even possible mass loss effects that act to slow the orbital evolution of the planet \\citep{Li10}.  In this paper, we present the discovery and characterisation of the hot Jupiter WTS-2 b. It is the second planet to be detected in the infrared light curves of the WFCAM Transit Survey (WTS) \\citep{Cap12,Kov12}, and orbits a mid-K dwarf star at just 1.5 times the separation at which it would be destroyed by tidal forces, making it a useful benchmark in constraining the theory of stellar tides. The remainder of this paper is organised as follows: in Section~\\ref{sec:WTS}, we briefly summarise the goals of the WTS and its atypical observing strategy, the reduction procedure used to generate the infrared light curves, and the processes used to identify WTS-2 b as a transiting candidate and the checks performed before proceeding with its follow-up observations. Section~\\ref{sec:obs} describes all of the follow-up data we obtained for WTS-2 b and their data reduction. We characterise the WTS-2 host star in Section~\\ref{sec:char_star}, and derive the corresponding properties of its planetary companion WTS-2 b in Section~\\ref{sec:char_planet}. Section~\\ref{sec:falsepos} summarises our investigation into possible blending scenarios. In Section~\\ref{sec:discussion}, we calculate and discuss the tidal evolution and remaining lifetime of WTS-2 b. We calculate the expected shift in its transit arrival time after 10 years, assuming that its orbit is decaying under tidal forces with $\\qprime=10^{6}$, and we give a correction to the previously published expected shift in the transit arrival time for WASP-18 b in Section~\\ref{sec:wasp-18}. In Section~\\ref{sec:pop}, we also attempt to constrain $\\qprime$ using the known population of hot Jupiters. Finally, we assess the potential for characterising the atmosphere of WTS-2 b using ground-based telescopes in Section~\\ref{sec:followup}. Our conclusions are summarised in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have reported the discovery of WTS-2 b, a typical transiting hot Jupiter in an unusually close orbit around a K2V star, which has a likely gravitationally-bound M-dwarf companion at a projected separation of $0.6$ arcsec. The proximity of the planet to its host star places it at just 1.5 times the separation at which it would be destroyed by Roche lobe overflow. The system provides a calibration point for theories describing the effect of tidal forces on the orbital evolution of giant planets, which are poorly constrained by observations. In particular, the system is useful for constraining theories that predict host stars with deeper convective envelopes lead to more efficient tidal dissipation. Using a simple model of tidal orbital evolution with a tidal dissipation quality factor $\\qprime=10^{6}$, we calculated a remaining lifetime for WTS-2 b of just $40$ Myr. The decaying orbit corresponds to a shift in the transit arrival time of WTS-2 b of $\\sim17$ seconds after $15$ years. We have also reported a correction to the previously published predicted shift in the transit arrival time of WASP-18 b, which used a very similar model for the stellar tides. We have calculated that the WASP-18 b transit time shift is $356$ seconds after 10 years for $\\qprime=10^{6}$, which is much larger than the previously reported $28$ seconds. We found that transit arrival time measurements in individual systems could place stringent observational constraints on $\\qprime$ across the full mass spectrum of exoplanet host stars within the next decade.  Our attempt to constraint $\\qprime$ via a study of the observed population of currently known transiting hot Jupiters was inconclusive, requiring a more detailed and precise determination of transit survey sensitivities at long periods. We conclude that the most informative and straight-forward constraints on $\\qprime$ and the theory of tidal orbital evolution for exoplanets will be provided by transit arrival time shifts in individual systems. Finally, WTS-2 b is one of the most highly irradiated gas giants orbiting a K-dwarf and is therefore expected to have an inversion layer in its atmosphere. This is in contrast to the non-inverted atmosphere of HD 189733 b, which has a very similar planet mass and host star to WTS-2 b, but receives $\\sim3$ times less incident radiation. Despite the relatively faint magnitude of the host star, the system size ratio and hot day-side temperature result in predicted infrared secondary eclipses that are within the reach of current ground-based instrumentation."
    },
    "1402/1402.0005_arXiv.txt": {
        "abstract": "We adopt the conduction fluid approximation to model the steady-state distribution of matter around a massive black hole at the center of a weakly collisional cluster of particles. By ``weakly collisional\" we mean a cluster in which the mean free time between particle collisions is much longer than the characteristic particle crossing (dynamical) time scale, but shorter than the cluster lifetime.  When applied to a star cluster, we reproduce the familiar Bahcall-Wolf power-law cusp solution for the stars bound to the black hole. Here the star density scales with radius as $r^{-7/4}$ and the velocity dispersion as $r^{-1/2}$ throughout most of the gravitational well of the black hole. When applied to a relaxed, self-interacting dark matter (SIDM) halo with a velocity-dependent cross section $\\sigma \\sim v^{-a}$, the gas again forms a power-law cusp, but now the SIDM density scales as $r^{-\\beta}$, where $\\beta = (a+3)/4$, while its velocity dispersion again varies as $r^{-1/2}$. Results are obtained first in Newtonian theory and then in full general relativity. Although the conduction fluid model is a simplification, it provides a reasonable first approximation to the matter profiles and is much easier to implement than a full Fokker-Planck treatment or an $N$-body simulation of the Boltzmann equation with collisional perturbations. ",
        "introduction": "Determining the stellar distribution around a massive black hole at the center of a virialized star cluster is an interesting and  well-studied problem. Originally formulated by Peebles~\\cite{Pee72a, Pee72b}, the problem was solved for the profile in a spherical globular cluster by Bahcall and Wolf~\\cite{BahW76}. They assumed an isotropic velocity distribution function to solve numerically the one-dimensional Fokker-Planck equation for the steady-state stellar distribution function $f(E)$ of stars bound to the black hole. Here $E$ is the energy per unit mass of a star. They found that the cumulative, distant, two-body encounters (i.e. small-angle Coulomb scattering) between stars drives the density throughout most of the gravitational well of the black hole (the ``cusp'') to a power-law profile that varies with radius as $\\rho \\propto r^{-7/4}$, while the velocity dispersion varies as $v \\propto r^{-1/2}$. The effects of velocity anisotropy and the role of the stellar disruption loss cone in the cusp were delineated by Frank and Rees~\\cite{FraR76} and Lightman and Shapiro~\\cite{LigS77}. They showed that while the distribution function depends only logarithmically on $J$, the angular momentum per unit mass of a star, the inward flux of stars and stellar disruption rate by the black hole are affected more significantly. They also applied the results to cusps in dense galactic nuclei as well as star clusters. All of this work motivated detailed two-dimensional Fokker-Planck treatments to determine $f(E,J)$, both by Monte Carlo~\\cite{ShaM78, MarS79} and by finite-difference~\\cite{CohK78} methods. The role of the removal of bound stars by the black hole in heating the ambient star cluster and reversing secular core collapse (i.e. reversing the ``gravothermal catastrophe'') was pointed out by Shapiro~\\cite{Sha77} and studied numerically via time-dependent Monte Carlo simulations of the two-dimensional Fokker-Planck equation to determine  $f(E,J;t)$~\\cite{MarS80}. (For a review of some of this early work and references see~\\cite{Sha85}). The same problem has received considerable attention in recent years by many authors, focusing on such aspects as nonspherical clusters [e.g.,~\\cite{VasM13}, mass segregation~\\cite{AleH09}, resonant relaxation~\\cite{RauT96}, relativistic corrections, tidal disruptions, and extreme-mass ratio inspirals (EMRIs) (see~\\cite{Ama12} for a review and references)].  In this paper we return to the the steady-state distribution of matter around a massive black hole at the center of a virialized, weakly collisional, spherical gas, such as a star cluster. Here we show that the conduction fluid model provides a straightforward means of deriving the steady-state matter density and velocity dispersion profiles in the cluster. It is approximate in that it treats the lowest order moments of the distribution function, and not the distribution function itself, and adopts several physically reasonable, but simplifying, relations to close the moment equations. However, the approach is much easier to implement than a full Fokker-Planck treatment or an $N$-body simulation of the Boltzmann equation with collisional perturbations. We then use this approach to determine the matter distribution in a cusp formed well inside the homogeneous core of a self-interacting dark matter (SIDM) halo containing a massive, central black hole.  SIDM with cross sections per unit mass in the range $\\sigma/m \\approx 0.45-450 ~{\\mbox{cm}^2\\;\\mbox{gm}^{-1}}$ [or $8 \\times 10^{-(25-22)} ~{\\mbox{cm}^2\\;\\mbox{GeV}^{-1}}$] was proposed by Spergel and Steinhardt \\cite{SpeS00} to rectify the problem that cosmological simulations~(e.g., \\cite{NavFW97, BulKSSKKPD01, WecBPKD02}) of purely collisionless cold dark matter (CDM) exhibit dark matter halos that are highly concentrated and have density cusps, in contrast to the constant density cores observed for galaxies and galaxy clusters.  The SIDM remedy fell out of favor for various reasons, including possible alternative explanations for this apparent density discrepancy (e.g. feedback and expulsion of baryonic gas), Bullet cluster observations that seemed to constrain SIDM models to have a cross section too small to have a significant effect on the structure of dark matter halos, gravitational lensing and x-ray data suggesting that the cores of galaxy clusters are denser and less spherical than predicted for SIDM, as well as theoretical biases. However, the interpretation of some of the data is not definitive (see \\cite{RocPBKGOM13} for a discussion and references), while new observational discrepancies with dissipationless CDM models have become apparent, such as the absence of dwarf spheroidal galaxies (dSphs) predicted by CDM simulations (the ``missing satellite problem''). Moreover, recent SIDM simulations~{\\cite{RocPBKGOM13, PetRBK13} now suggest that SIDM systems with a (velocity-independent) cross section as large as $\\sigma/m  \\approx 0.1 ~{\\mbox{cm}^2\\;\\mbox{gm}^{-1}}$ can yield the core sizes and constant central densities observed in dark matter halos at all scales, from clusters, to low surface brightness spirals (LSBs) and dSphs, and are consistent with revised Bullet cluster observational constraints. All of the other large-scale triumphs of CDM appear to be matched by these SIDM simulations. In addition, simulations involving velocity-dependent SIDM cross sections also appear to be successful in forming small, less concentrated cores and altering the inner density structure of subhalos in a way that is compatible with observations of dSphs but having a negligible effect on galaxy cluster scales~{\\cite{FirDACH00, FenKTY09, BucF10, LoeW11, VogZL12}. Such velocity-dependent interactions arise in ``hidden'' sector extensions to the Standard Model that have been constructed to explain some charged-particle cosmic ray observations in terms of dark matter annihilations~{\\cite{PosRV08, ArkFSW09, FoxP09, FenKY10a}.  The renewed viability of SIDM models motivates our own interest in returning to a subject that we studied earlier. We previously explored the nature of the ``gravothermal catastrophe'' (secular core collapse) in isolated SIDM halos~\\cite{BalSI02} using the conduction fluid formalism (see also~\\cite{KodS11} and references therein) and then showed that the formation of a black hole in the center of a galactic halo may be a natural and inevitable consequence of gravothermal evolution~\\cite{BalS02}.  Supermassive black holes (SMBHs) with masses in the range  $10^6-10^{10} M_{\\odot}$ are believed to be the engines that power active galactic nuclei and quasars. There is also substantial evidence that SMBHs reside at the centers of many, and perhaps most, galaxies~\\cite{Ric98, Ho99}, including the Milky Way~\\cite{GenEOE97, EckGOS02, GheSHTLMBD05}. While SMBHs may arise from baryonic processes, such as the collapse of Pop III stars~\\cite{MadR01} or supermassive stars~\\cite{ShiS02}, followed by mergers and gas accretion, their origin is not yet known.  It is therefore not unreasonable to imagine that bound structures containing dark matter of {\\it all} sizes (galaxy clusters, galaxies, satellite systems, etc) contain massive, central black holes. This possibility motivates the study in this paper, for which we adopt a velocity-dependent SIDM cross section $\\sigma \\propto v^{-a}$ and consider a relaxed, spherical, SIDM core with a massive black hole at its center and determine the cusp profile well inside the core. We note that a cusp can also form in a purely collisionless CDM system containing a central black hole. For example, if a central black hole grows adiabatically (e.g. by gas accretion) it will perturb the CDM particle orbits while holding the adiabatic invariants of the motion constant, and this process will result in a cusp~\\cite{Pee72a, GonS99, SadFW13}. For a SIDM system, however, particle collisions will wash out any initial and/or adiabatically altered particle distribution in a collisional relaxation time scale and the cusp will relax to the solution determined below.  In Sec. II we summarize our basic model and the assumptions that define it. In Sec. III we present our formulation of the problem in Newtonian theory and use simple scaling to derive the power-law cusp solution for bound particles obeying a power-law velocity-dependent interaction cross section (SIDM matter or stars). We also calculate the energy flux conducted outward by the particles bound to the black hole in the cusp and discuss how this flux may eventually halt and reverse secular core collapse. In Sec. IV we reformulate the problem in general relativity and integrate the resulting equations numerically to obtain the full density and velocity profiles in the cusp. In Sec. V we discuss our findings, evaluate the contribution of unbound particles, and consider the very different behavior expected at very late times when a SIDM core evolves to a collisional fluid. We also summarize some future work.  We adopt geometrized units and set $G=1=c$ below. ",
        "conclusions": "The case $a=4$ $(\\sigma \\propto v^{-4})$ is particularly interesting, both for star clusters and SIDM halos. This value corresponds to a Coulomb scattering cross section and therefore yields the familiar cusp profiles for a star cluster containing a massive, central black hole: $\\rho_N \\propto r^{-7/4}$ and $v \\propto r^{-1/2}$. Our numerical results are in close agreement with the steady-state profile found by Bahcall and Wolf~\\cite{BahW76}, who integrated the Fokker-Planck equation for a star cluster satisfying the same assumptions adopted here (see Sec.~\\ref{sec:assump}). The analytic fitting formulas defined in Eqs.~(\\ref{eq:analdens}) and (\\ref{eq:analvel}), with the very coefficients for these formulas quoted above for $a=4$, were in fact constructed by Bahcall and Wolf to fit their own numerical solution for star clusters. It is reassuring to discover that these formulas also provide good fits to our numerical solution for this case, based as it is on a much simpler formulation and calculation.  Interestingly, the same $a=4$ case is also relevant for some new models proposed for SIDM systems. A velocity-dependent, non-power-law, elastic scattering cross section, motivated by a Yukawa-like potential involving a new  gauge boson of mass $m_{\\phi}$, (and similar to ``screened'' Coulomb scattering in a plasma) has been proposed recently~\\cite{FenKY10b, FinGSVW11}. $N$-body simulations adopting this interaction apparently can explain the observed cores in dwarf galaxies without affecting the dynamics of larger systems with larger velocity dispersions, such as clusters of galaxies~\\cite{LoeW11, VogZL12}.  Now the presence of a massive, central black hole in the core of a SIDM system will develop a cusp, as we have discussed. For any interaction cross-section, power law or not, the density and velocity profiles in such a cusp would be straightforward to calculate by the method formulated in this paper. However, we can already anticipate the results for systems interacting via the above Yukawa-like potential, provided their core velocity dispersions $v$ satisfy $\\beta_c \\equiv \\pi (v_{\\rm max}/v)^2 \\lesssim 0.1$. In this case the cross sections vary with velocity according to $\\sigma \\sim \\beta_c^2 \\sim v^{-4}$, up to a slowly varying logarithmic factor. Here $v_{\\rm max}$ is the velocity at which the momentum-transfer weighted scattering rate $\\langle\\sigma v\\rangle$ peaks at a cross section value of $\\sigma^{\\rm max} = 22.7/m_{\\phi}^2$. When $v_{\\rm max}$ is comparable to the velocity dispersion in the core of a typical dwarf spheroidal galaxy, the $N$-body simulations that result are reported to give profiles in better agreement with observations for these objects than collisionless CDM predictions. Adopting such a value for $v_{\\rm max}$ implies that throughout the bulk of a cusp in dwarf galaxies, as well as in any larger system, the interaction cross section resides in the $n=4$ power-law regime, for which the Bahcall-Wolf solution applies in a first approximation.  \\subsection{Unbound particles}  In addition to the bound particles that orbit in the cusp there are {\\it unbound} particles ($E/m > 1$, where $E$ again includes rest-mass energy) from the ambient core that penetrate into the cusp. Treating the core as an infinite, collisionless and monoenergetic bath of such particles with rest-mass density $\\rho_0$ moving isotropically with (one-dimensional) velocity dispersion $v_0$ yields a density profile in the gravitational well of the black hole given by \\begin{equation} \\label{eq:unbound} \\rho_N/\\rho_0 = (1 + 2r_h/3r)^{1/2}, \\end{equation} and a velocity profile \\begin{equation} v/v_0 = (1 + 2r_h/3r)^{1/2}, \\end{equation} assuming Newtonian gravitation~\\cite{ZelN71, ShaT83} (for the solution in general relativity, see \\cite{ShaT85}). The key point is that a comparison of Eqs.~(\\ref{eq:powlaw}) or (\\ref{eq:analdens}) with (\\ref{eq:unbound}) shows that the density of unbound particles penetrating the interior of the cusp is much smaller than the density that builds up in bound particles for all plausible ($a>0$) velocity-dependent interactions.   \\subsection{Collisional fluid regime}  As shown in~\\cite{BalSI02}, gravothermal evolution in a SIDM cluster characterized by a velocity-{\\it independent} self-interaction inevitably drives the contracting core to sufficiently high density that the particle mean free path to collisions becomes smaller than the local scale height in the innermost regions. A similar situation presumably arises for a range of velocity-dependent interactions. When this situation occurs, the particles in the innermost regions will behave as a collisional (hydrodynamical) fluid. A massive black hole at the center of such a core would then be expected to accrete particles via steady-state, adiabatic, spherical Bondi~\\cite{Bon52} flow. The cusp in such a case will fill with gas that has a density profile of the form \\begin{equation} \\label{eq:densbondi} \\rho_N/\\rho_0 \\approx 1 + \\chi (r_h/r)^{3/2},  \\ \\ \\ \\chi \\sim \\mathcal{O}(1), \\end{equation} and an inward radial velocity approaching the free-fall velocity $v_{ff} \\sim (M/r)^{1/2}$ inside $r_h$. The temperature, or velocity dispersion of the particles, $ \\gamma v^2 \\sim k_B T/m$, will satisfy \\begin{equation} \\label{eq:tempbondi} T/T_0 \\approx (\\rho_N/\\rho_0)^{\\Gamma-1}. \\end{equation} Here the adiabatic index $\\Gamma$ is 5/3 for a nonrelativistic gas ($k_B T/m \\ll 1$) and $4/3$ for a relativistic gas ($k_B T/m \\gg 1$) and varies with radius from the black hole (see~\\cite{ShaT83} for a review of spherical Bondi flow, including a relativistic treatment). For a nonrelativistic core temperature $T_0$, $\\Gamma$ will be near $5/3$ throughout most of the cusp, decreasing below $5/3$  as the fluid approaches the horizon. In this case Eqs.~(\\ref{eq:densbondi}) and (\\ref{eq:tempbondi}) show that the temperature of the fluid at the horizon will be marginally relativistic, climbing to  $k_BT_{\\rm horz}/m \\sim 1$, independent of $T_0$.  The ratio of the collision mean free path to the characteristic scale height in the core of a typical SIDM model constructed to explain currently observed clusters satisfies \\begin{eqnarray} \\label{eq:mfp} \\frac{\\lambda_0}{R_0} &=& \\frac{1}{\\rho_0 \\sigma_0 R_0} \\cr &\\approx& 50 \\left(\\frac{0.1 ~M_{\\odot}/{\\rm pc}^3}{\\rho_0}\\right) \\left(\\frac{1 ~{\\rm cm}^2/{\\rm g}}{\\sigma_0}\\right) \\left(\\frac{1 ~{\\rm kpc}}{R_0}\\right), \\end{eqnarray} and this ratio increases for $r \\gtrsim R_0$. The above inequality is consistent with, e.g., the velocity-dependent SIDM models constructed in \\cite{VogZL12} (see their Fig. 4 for main haloes and Fig. 7 for subhaloes). Inside a cusp around a central black hole the ratio becomes \\begin{equation} \\frac{\\lambda}{r} = \\frac{1}{\\rho \\sigma r} \\approx \\frac{\\lambda_0}{R_0} \\left(\\frac{M_0}{M}\\right) \\left(\\frac{r_h}{r}\\right)^{(1+a)/4}, \\ \\ \\ r \\lesssim r_h, \\end{equation} where we have used Eq.~(\\ref{eq:powlaw}), and thus increases as $r$ decreases below $r_h$ for all $a > -1$. Thus a typical SIDM cluster with a central black hole is characterized by $\\lambda/r \\gg 1$ everywhere and the matter resides in the weakly, and not the fluid, collisional regime.  Now SIDM clusters {\\it can} evolve from a weakly collisional gas to a strongly collisional fluid, transforming the nature of the cusp profiles accordingly. However, SIDM models constructed to explain observed clusters are specifically designed {\\it not} to have undergone such evolution as yet. In particular, these models employ interaction cross sections whose magnitudes are sufficiently small that secular core collapse -- the gravothermal catastrophe -- will not have occurred in a Hubble time. In these cases there is insufficient time for the cores to have evolved to a fluid state, which requires a significant fraction of a core collapse time scale ($\\approx 290 t_r$ for isolated clusters with a=0; \\cite{BalSI02}). This conclusion is evident from the relation $t_r/t_{\\rm dyn} \\approx \\lambda_0/R_0$, where $t_{\\rm dyn} \\approx 1/\\rho_0^{1/2}$ is the characteristic dynamical time scale in the core. We thus have $t_r \\approx 5\\times10^7 (\\lambda_0/R_0)(0.1 ~M_{\\odot} {\\rm pc}^{-3}/ \\rho_0)^{1/2}$ yr, whereby substituting Eq.~\\ref{eq:mfp}, yields a typical core collapse time scale significantly longer than the Hubble time. Moreover, core collapse is suppressed in (nonisolated) clusters when merging from hierarchical formation is included~\\cite{AhnS05}.   \\subsection{Future work}  The possibility of a massive black hole at the centers of SIDM clusters and the formation of a cusp around the black hole motivates several interesting questions that we are now investigating.  One question is: assuming the SIDM particles experience weak (WIMP-like) inelastic interactions in addition to the elastic interactions assumed here, what are the observable signatures of such a cusp?  Will, for example, the perturbative inelastic interactions produce detectable radiation or high-energy particles above the values expected from a homogeneous SIDM core without a cusp~\\cite{FiePS13}?  Another question is, given that the cusp serves as a source of kinetic energy conducted into the ambient core, is the energy sufficient to reverse secular core collapse and cause the cluster to reexpand, as discussed in Sec.~\\ref{sec:energy_flux}~\\cite{BalS13}?  These are just some of the issues that we will pursue in future studies of SIDM clusters.  {\\it Acknowledgments}: It is a pleasure to thank S. Balberg, B. Fields, A. Peter and P. Shapiro for useful discussions.  This paper was supported in part by NSF Grants No. PHY-0963136 and No. PHY-1300903 and NASA Grants No. NNX11AE11G and No. NN13AH44G at the University of Illinois at Urbana-Champaign. V. P. gratefully acknowledges support from a Fortner Fellowship at UIUC."
    },
    "1402/1402.7139_arXiv.txt": {
        "abstract": "Particle acceleration in magnetic reconnection of electron-positron plasmas is studied by using a particle-in-cell simulation. It is found that significantly large number of non-thermal particles are generated by the inductive electric fields around an X-type neutral line when the reconnection outflow velocity, which is known to be an \\Alfven velocity, is of the order of the speed of light. In such a relativistic reconnection regime, we also find that electrons and positrons form a power-law-like energy distribution through their drift along the reconnection electric field under the relativistic Speiser motion. A brief discussion of the relevance of these results to the current sheet structure, which has an anti-parallel magnetic field in astrophysical sources of synchrotron radiation, is presented. ",
        "introduction": "The non-thermal particles in space plasmas have attracted our attention for a long time. Their origins remain central problems and many kinds of acceleration processes have been studied. Double layer and shock acceleration mechanism are important ones.  Magnetic reconnection is one of the fundamental processes in plasmas. It is well accepted that it plays a crucial role in the Earth's magnetotail, solar corona and astronomical accretion disks. Since the stored magnetic energy is rapidly released to particle kinetic energy, a reconnection process is also important as one of the acceleration mechanisms in plasmas. In astronomical high-energy situations, several radio observations \\citep{lesch92,dejager94} have suggested that magnetic reconnection may be the source of non-thermal radiations. Magnetic reconnection in pair-plasmas is also discussed for energetics of a striped wind of the Crab pulsar \\citep{michel82,coro90,michel94,lyu01}.  According to the studies on reconnection of the Earth's magnetotail, particle acceleration in reconnection mainly takes place around the X-type neutral region, where the amplitude of magnetic fields is weak and charged particles become unmagnetized. They are accelerated by an inductive electric field perpendicular to the two-dimensional reconnection magnetic fields. Based on the studies on particle behavior around the X-type region \\citep{Speiser65,Sonnerup}, many authors \\citep{zele90,Deeg91} have investigated the particle energetization by test particle simulations, in the time-stationary/dependent electric and magnetic fields obtained by MHD simulations. They have demonstrated that accelerated particles form power-law-like energy spectra.  In order to study the acceleration in reconnection fields more precisely, we should consider self-consistency between particle motion and electromagnetic fields. Thus a full-particle simulation that follows the kinetic plasma equations is required. However, many of the full-particle studies of reconnection have focused on field structure or energy conversion. The acceleration of energetic particles in reconnection is not discussed enough in spite of its importance. \\citet{hoshi_supra} have discussed supra-thermal electrons accelerated near the magnetic pile-up region, in addition to the X-type region.  \\begin{figure*}[thp] \\plotone{f1.ps} \\caption{A snapshot of magnetic reconnection region of $-25.6 \\le X/\\lambda \\le 25.6$ and $-9.6 \\le Z/\\lambda \\le 9.6$ at $t / \\tau_A = 80.6$. The solid lines and the vectors show magnetic field lines and the plasma flows, respectively. The color contour represents the plasma density, which is normalized by the initial density in the plasma sheet $n_0$. \\label{rec}} \\end{figure*}   In this letter, we study particle acceleration in magnetic reconnection of astronomical pair-plasmas, using PIC (particle-in-cell) code. We chose the condition that typical \\Alfven velocity of plasma is the order of light speed. From our simulation results, we have found remarkable amount of non-thermal component in the energy spectrum and found that this is a sign of acceleration by the strong inductive electric field near the X-type region. The high-energy particles are accelerated through Speiser/meandering-like orbit around there, and their maximum energy has grown up to $mc^2\\Omega_cT$, where $\\Omega_c$ is the cyclotron frequency, and $T$ is the typical reconnection time which may be defined by $L/c$, where $L$ is the size of the reconnection region. Moreover, the spectra around the X-type region seem to have a power-law distribution with the power-law index of 1. We consider that feedback effect of relativistic inertia plays an important role in formation of the spectra, and propose a new process of formation of this spectrum.  This letter is organized as follows. In the next section we present our simulation conditions. In Sec. 3, we show several results and findings of our run. In Sec. 4, we study the acceleration in the vicinity of the X-type neutral point and introduce basic idea of formation of power-law spectrum. In Sec. 5, we summarize and conclude our study. ",
        "conclusions": "We have carried out simulations of magnetic reconnection in astronomical pair-plasmas and we found that significantly large number of non-thermal particles are generated due to the strong inductive electric field in the AR. Our simulation shows that particle acceleration by magnetic reconnection can work effectively in the universe. Moreover, the energy spectrum around the AR resembles power-law distribution with the power-law index of 1. This is a quite hard spectrum, which we have never yet seen.  It is highly expected that reconnection in relativistic pair-plasmas may occur in high-energy astronomical objects, such as a striped wind of pulsars and Active Galactic Nuclei. The acceleration by reconnection occurs in relatively short time scale, several tens of the \\Alfven transit time $\\tau_A = \\lambda / V_{A} \\sim \\lambda / c$, this acceleration may be a good candidate of the origin of non-thermal particles.  In our simulation code, we neglect any Coulomb collisions for simplicity. This treatment is acceptable since particles' mean free path is generally quite long in space plasmas. They are virtually collision-less, thus pair-annihilation is also neglected. Considering the fact that reconnection occur in a relatively short time scale $\\sim 10^1 \\times \\lambda / c$, our assumptions may be acceptable. Note that this code does not include any radiation losses, such as the synchrotron and the inverse Compton losses. In their test-particle study, \\citet{schopper} have noted that the effect of synchrotron loss is not negligible in their high-energy situation ($10^4 \\sim 10^6 mc^2$). The effects of pair-production and pair-annihilation are also neglected in our study. More precise simulation should include them but it can not be achieved using particle-in-cell code.  We have also discussed the formation of power-law spectrum around the X-type region and our simulation result is well described by $N \\propto \\varepsilon ^{-s}$, where the index $s$ is about the unity. Strictly speaking, it is unable to discuss the classical Speiser/meandering orbit in the vicinity of the X-type region where $|E_y| \\ge |B_{total}|$, since the characteristic of Lorentz transformation changes. Particle orbits are still similar to Speiser/meandering orbit, but they are driven by $E_y$ and resonant with $E_y$. We are still in the process of searching analytical solutions for the acceleration orbit, but the essential point of our model is as follows: the more strongly particles are accelerated in the Y direction, the longer time they stay in the AR, due to the relativistic effect of cyclotron-motion.  In our run, the growth of the maximum energy $\\varepsilon_{max}$ seems to be confined by the periodic boundary. Simulation run with larger simulation box may produce more energy than $38mc^2$. The multiple reconnections may be also important for producing high-energy particles. If two or more reconnection regions are formed in the plasma sheet, ejected particles travel into another AR so that they can be more accelerated.  We have also performed simulation runs with different \\Alfven speed parameters. We have recognized remarkable number of non-thermal particles in every case. Together with dependency of reconnection behavior on the plasma sheet thickness, we will report the full story in the coming paper."
    },
    "1402/1402.4006_arXiv.txt": {
        "abstract": "{Until recently, dust emission has been detected in very few host galaxies of gamma-ray bursts (GRBHs). With \\hers, we have now observed 17 GRBHs up to redshift $z\\sim3$ and detected seven of them at infrared (IR) wavelengths. This relatively high detection rate (41\\%) may be due to the composition of our sample which at a median redshift of 1.1 is dominated by the hosts of dark GRBs. Although the numbers are small, statistics suggest that dark GRBs are more likely to be detected in the IR than their optically-bright counterparts. Combining our IR data with optical, near-infrared, and radio data from our own datasets and from the literature, we have constructed spectral energy distributions (SEDs) which span up to 6 orders of magnitude in wavelength. By fitting the SEDs, we have obtained stellar masses, dust masses, star-formation rate (SFR), and extinctions for our sample galaxies. We find that GRBHs are galaxies that tend to have a high specfic SFR (sSFR), and like other star-forming galaxies, their ratios of dust-to-stellar mass are well correlated with sSFR. Dust masses of GRBHs relative to stellar mass and SFR fall within the range of other star-forming galaxies in the local universe, and of sub-millimeter galaxies (SMGs) and luminous IR galaxies for redshift $z\\ga1$. We incorporate our \\hers\\ sample into a larger compilation of GRBHs, after checking for consistency in mass and SFR estimations. This combined sample is compared to SFR-weighted median stellar masses of the widest, deepest galaxy survey to date in order to establish whether or not GRBs can be used as an unbiased tracer of cosmic comoving SFR density (SFRD) in the universe. In contrast with previous results, this comparison shows that GRBHs are medium-sized galaxies with relatively high sSFRs, as might be expected for galaxies selected on the basis of SFR because of the explosive GRB event. Stellar masses and sSFRs of GRBHs as a function of redshift are similar to what is expected for star-forming galaxy populations at similar redshifts. We conclude that there is no strong evidence that GRBs are biased tracers of SFRD; thus they should be able to reliably probe the SFRD to early epochs. ",
        "introduction": "\\label{sec:intro} Long-duration gamma-ray bursts (GRBs) are so luminous that they can shine through highly obscured galaxies \\citep[e.g.,][]{djor01} and can be seen even at very high redshifts \\citep{salvaterra09,tanvir09,cucchiara11}. They are thought to originate in the collapse of very massive stars at the end of their evolution \\citep{pacz98,macfayden99,woosley06}. Because of this association with massive stars, GRBs have recently been used, thanks to the advent of the dedicated mission {\\it Swift}, to infer the cosmic evolution of the star formation rate density (SFRD) up to $z \\sim 9$ \\citep{yuksel08,kistler09,butler10,robertson12,elliott12,trenti13}.  Although GRBs are rare events, they enable identification of galaxies that would not otherwise be singled out even in deep flux-limited surveys, thus making GRBs a potentially powerful probe of galaxy evolution. Galaxies hosting GRBs (GRBHs) are better known in the low-$z$ regime ($z\\la 1.5$), where they are typically low-mass, young, star-forming blue galaxies with low dust extinction \\citep{lefloch03,christensen04,fruchter06,savaglio09}. However, these galaxy characteristics, possibly related to a selection bias because of the consideration of only optically bright GRB afterglows, may not be so uniform at high redshift. Indeed, evidence is mounting that the GRBH population is much more diverse at $z\\ga 1.5$ than previously thought.  Dark GRBs, those for which the observed optical afterglow is very faint relative to the extrapolation from the X-ray \\citep{jakobsson04,vanderhorst09}, tend to be found in massive, star-forming galaxies with red colors, high extinction and large SFRs \\citep[e.g.,][]{kruhler11,rossi12,perley13}. Dark GRBs comprise up to 30-40\\% of the {\\it Swift} GRB dataset \\citep{fynbo09,greiner11,melandri12}, and thus the assumption that all GRBHs are low-mass, metal poor galaxies may be an oversimplification. Consequently, theoretical work based on this assumption using GRBHs as cosmological probes \\citep{campisi09,niino11,robertson12} could be undermined.  Until recently, statistics on dark GRBs and their host galaxies have been poor \\citep[e.g.,][]{kann10}. Because of their optical faintness, it has been very difficult to localize the optical afterglow of dark GRBs, and thus identify the host \\citep[e.g.,][]{rossi12}. Now, thanks to sustained observational efforts, we have a considerably better understanding of dark GRBs and their host galaxies. In most dark GRBs, the optical faintness is caused by high dust extinction columns and moderate redshift \\citep{perley09,greiner11,kruhler11,rossi12,melandri12,covino13}. However, the properties of dust extinction are not the same for all GRBs, although there is no clear evidence that afterglow extinction curves differ significantly from those commonly used \\citep{perley08,zafar11,zafar12,schady12}. It is thus difficult to make conclusive statements about the nature of GRB hosts \\citep[e.g.,][]{delia13,elliott13}; it is neither true that all massive, metal-rich GRBHs are found from dark GRBs \\citep[e.g.,][]{lefloch02}, nor do all dark GRBs reside in massive hosts at high redshift.  \\begin{center} \\begin{table*} \\caption[]{Host galaxy sample} \\label{tab:sample} \\resizebox{\\linewidth}{!}{ {% \\tiny \\begin{tabular}{llllccll} \\hline \\multicolumn{1}{c}{GRB} & \\multicolumn{1}{c}{RA} & \\multicolumn{1}{c}{Dec.} & \\multicolumn{1}{l}{Redshift} & \\multicolumn{1}{c}{Redshift} & \\multicolumn{1}{c}{Dark?$^{\\mathrm a}$} & \\multicolumn{2}{c}{Observing ID} \\\\ & \\multicolumn{2}{c}{(J2000)$^{\\mathrm{b}}$} & & \\multicolumn{1}{l}{Reference\\,$^{\\mathrm{c}}$} && \\multicolumn{1}{c}{PACS} & \\multicolumn{1}{c}{SPIRE} \\\\ \\hline \\\\ 970828                & 18:08:31.56 & +59:18:51.2   &  0.958  & 1 & Yes &  1342258609,10 & 1342241152 \\\\ % 980613                & 10:17:57.9  & +71:27:26.3   &  1.097  & 2 & No  &  1342270864,65 & \\multicolumn{1}{c}{$-$} \\\\ % 980703                & 23:59:06.67 & +08:35:06.7   &  0.966  & 3 & No  &  \\multicolumn{1}{c}{$-$}           & 1342212419 \\\\ % 990705                & 05:09:54.50 & $-$72:07:53.0   &  0.842  & 4 & No  &  1342261814,15 & 1342214749 \\\\ % 020127                & 08:15:01.42 & +36:46:33.9   &  1.9$^{\\mathrm d}$    & 5 & Yes &  1342243296,97 & 1342251906 \\\\ % 020819B               & 23:27:19.52 & +06:15:53.2   &  0.411  & 6 & Yes &  1342246708,09 & 1342212294 \\\\ % 030115                & 11:18:32.63 & +15:02:59.9   &   2.0$^{\\mathrm e}$   & 7  & Yes & 1342247632,33 & \\multicolumn{1}{c}{$-$} \\\\ % 050219A               & 11:05:39.07 & $-$40:41:04.6   &  0.2115  & 8 & Yes  & 1342248286,87 & 1342247261 \\\\ % 050223                & 18:05:32.99 & $-$62:28:18.8   &  0.584  & 9  & Yes & 1342243788,89 & 1342214756 \\\\ % 051022                & 23:56:04.1  & +19:36:24.1   &  0.807  & 6 & Yes  & 1342238085,86 & 1342247985 \\\\ % 060904A               & 15:50:54.56 & +44:59:10.5   &  2.55$^{\\mathrm d}$  & 11 & Yes & 1342261849,50 & \\multicolumn{1}{c}{$-$} \\\\ % 070306                & 09:52:23.3  & +10:28:55.5   &  1.496  & 12 & Yes & 1342254142,43$^{\\mathrm{f}}$ & \\multicolumn{1}{c}{$-$} \\\\ % 071021                & 22:42:34.31 & +23:43:06.5   &  2.452  & 13 & Yes & 1342246178,79 & 1342258356 \\\\ % 080207                & 13:50:03.01 & +07:30:07.8   &  2.086  & 13 & Yes & 1342257547,48 & 1342261526 \\\\ % 080325                & 18:31:34.3  & +36:31:24.8   &  1.78   & 14 & Yes & 1342245672,73 & 1342241156 \\\\ % 090404                & 15:56:57.52 & +35:30:57.5   &  3.00$^{\\mathrm d}$  & 14 & Yes & 1342258433,34 & 1342241163 \\\\ % 090417B               & 13:58:46.66 & +47:01:04.4   &  0.345  & 15 & Yes & 1342257593,94 & 1342259467 \\\\ % \\\\ \\hline \\end{tabular} } } \\vspace{0.5\\baselineskip} \\begin{description} \\item [$^{\\mathrm{a}}$] These GRBs are classified as dark according to the definitions by \\citet[][]{jakobsson04,vanderhorst09}. \\item [$^{\\mathrm{b}}$] These are the positions of the host galaxy, which may or may not be exactly coincident with the position of the GRB afterglow. \\item [$^{\\mathrm{c}}$] (1) \\citet{djor01}; (2) \\citet{djor03}; (3) \\citet{djor98}; (4) \\citet{lefloch02}; (5) \\citet{berger07}; (6) \\citet{levesque10b}; (7) \\citet{levan06}; (8) \\citet{rossi14}; (9) \\citet{pellizza06}; (10) \\citet{chary07}; (11) \\citet{xiao11}; (12) \\citet{jaunsen08}; (13) \\citet{kruhler12}; (14) \\citet{perley13}; (15) \\citet{holland10}. \\item [$^{\\mathrm{d}}$] Photometric redshift. \\item [$^{\\mathrm{e}}$] Photometric redshift determined by our \\grasil\\ fits. \\item [$^{\\mathrm{f}}$] Although this source was in our target list, this observation was acquired in the OT2 proposal, OT2\\_ppschady\\_2. \\end{description} \\end{table*} \\end{center}   Direct observations of the entire spectral energy distribution (SED) of dark GRBHs would help to place them in the context of other high-$z$ galaxy populations, especially since the most common understanding of the high-$z$ universe is based on optically-selected (rest-frame UV) galaxy surveys. So far, only a handful of GRBHs have been detected with sub-millimeter (submm) facilities \\citep{barnard03,berger03,tanvir04,priddey06,wang12,michalowski13}. However, in this minority SFRs can be high, $\\sim$500\\,\\sfr, as high as those of the submm galaxy (SMG) population \\citep{chapman04,greve05} even though the galaxies picked out by the two selection criteria are quite different. There has been no CO emission found in any of the GRBHs observed so far \\citep[e.g.,][]{hatsukade11} and there is some hint that dust in GRBHs may be warmer than in typical ultra-luminous infra-red galaxies (ULIRGs) and SMGs \\citep[][]{priddey06,michalowski08}.  To better understand the properties of dark GRB hosts, and to assess their potential impact on the GRBH population in general, we undertook an observational campaign with \\hers; this paper presents results from this campaign. We characterize the dust content, stellar mass, and SFRs of GRBHs through \\hers\\ observations of 17 GRBHs, 14 of which host dark GRBs. This is the first time that \\hers\\ has been used to examine dust in GRBHs; dust emission is detected in seven of our targets. Sample selection is described in Sect. \\ref{sec:sample}, and Sect. \\ref{sec:data} reports the \\hers\\ observations and the other data incorporated in the compilation of the SEDs, together with the procedures for the photometry. SED fitting is discussed in Sect. \\ref{sec:grasil}, and Sect. \\ref{sec:results} gives the results of the fitting in terms of stellar masses, dust masses, SFRs, and dust extinction. The properties of our sample of GRBHs are compared with other GRBH samples and other high-$z$ star-forming galaxy populations in Sect. \\ref{sec:context}. Throughout the paper, we assume a $\\Omega_m$\\,=\\,0.3, $\\Omega_\\Lambda$\\,=\\,0.7 cosmology, with Hubble constant $H_0$\\,=\\,70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions}  For the first time, we are able to place constraints on dust emission in a significant sample of the hosts of long GRBs. We have observed with \\hers\\ 17 GRBHs from $z\\sim0.2$ to $z\\sim3$, and detected dust emission in 7. The probability of IR detection in the dark subset of our sample is 43\\%, compared to $\\la$20\\% from previous attempts to detect IR emission which targeted hosts of optically-bright GRBs \\citep[e.g.,][]{tanvir04,lefloch06}. Fitting the multiwavelength SEDs with \\grasil\\ gives dust and stellar masses, SFRs, and dust extinction, and enables a comparison of these properties in GRBHs with other star-forming galaxy populations at similar redshift. This comparison shows that GRBHs at $z\\ga$0.5 tend to be galaxies with high sSFRs and high dust-to-stellar mass ratios, but are in any case similar in terms of \\mstar, SFR, and \\av\\ to other populations. The \\mdust/\\mstar\\ ratios and sSFR of GRBHs are similar to SMGs and ULIRGs at similar redshifts. The trends of \\mdust/SFR suggest that GRBs may select warm SMG-like objects, which before \\hers\\ were difficult to identify.  To investigate whether GRBs % can be used to trace co-moving SFR in the universe, we have incorporated additional GRBHs in our analysis and compared them with trends of redshift for \\mstar\\ and sSFR as given by the recent UltraVISTA survey of 220\\,000 galaxies ($\\sim$135\\,000 star-forming). The results show that the stellar masses of GRBHs in our combined sample lie within the range of 50\\% of the UltraVISTA star-forming galaxy population from $z\\sim$0.3 to $z\\sim3$, although there is a large dispersion. The sSFRs of the GRBHs tend to be high, but are also within the range of values expected for star-forming galaxies at similar redshifts. Thus we conclude that GRBs select galaxies that are representative of the more general population; % hence they should trace the cosmic SFRD in an unbiased way.  Our sample is dominated by hosts of dark GRBs, and dark hosts have \\mstar\\ above the GRBH median in all redshift bins. Because dark and dusty bursts tend to be found in more mature, metal-rich galaxies, it is possible that this sample composition is driving our conclusion that GRBH are unbiased tracers of star formation. Better statistics are needed, together with careful SED fitting, to better address the properties of the hosts of dark GRBs, their impact on the GRBH population as a whole, and whether the hosts of GRBs are generally representative of high-$z$ galaxies."
    },
    "1402/1402.1623_arXiv.txt": {
        "abstract": "We describe a new iterative approach for the realization of equilibrium N-body systems for given density distributions. Our method uses elements of Schwarzschild's technique and of the made-to-measure method, but is based on a different principle. Starting with some initial assignment of particle velocities, the difference of the time-averaged density response produced by the particle orbits with respect to the initial density configuration is characterized through a merit function, and a stationary solution of the collisionless Boltzmann equation is found by minimizing this merit function directly by iteratively adjusting the initial velocities. Because the distribution function is in general not unique for a given density structure, we augment the merit function with additional constraints that single out a desired target solution.  The velocity adjustment is carried out with a stochastic process in which new velocities are randomly drawn from an approximate solution of the distribution function, but are kept only when they improve the fit. Our method converges rapidly and is flexible enough to allow the construction of solutions with third integrals of motion, including disk galaxies in which radial and vertical dispersions are different. A parallel code for the calculation of compound galaxy models with this new method is made publicly available. ",
        "introduction": "\\label{sec:intro}  The large number of stars and dark matter particles in galaxies and galaxy clusters makes them essentially perfect collisionless systems. Their dynamics is hence described by the collisionless Boltzmann equation, coupled to self-gravity through Poisson's equation. Relaxed systems correspond to stationary solutions of these equations, and much of the field of galactic dynamics is concerned with understanding different aspects of these solutions \\citep[see][for an excellent exposition]{BinneyTremaine2008}. This is particularly important for using observational probes of kinematics to infer, for example, something about the underlying density distribution.  Numerical N-body simulations have become a primary work-horse to study collisionless systems, both in stationary and dynamic situations. Prominent examples include the study of bar instabilities \\citep[e.g.][]{Athanassoula2002}, the formation of spiral waves \\citep[e.g.][]{Donghia2013}, or major and minor mergers of galaxies \\citep[e.g.][]{Barnes1992, Hernquist1995}.  They are also actively used to study the response of disks to the bombardment by dark matter clumps \\citep[e.g.][]{KazantzidisI2008, Donghia2010}, or the radial migration of stars caused by resonance scattering \\citep[e.g.][]{Sellwood2002}, and many more.   In carrying out numerical experiments targeting these questions, a recurrent challenge is to construct suitable initial conditions.  One usually requires them to be in a reasonably stable, approximate equilibrium in the beginning, otherwise any subsequent dynamics may be dominated or heavily contaminated by the specific out-of-equilibrium state one started out with. Often, one has a relatively clear notion of the density structure one wants to realize, but initializing the particle velocities appropriately is a quite non-trivial problem. This is because doing this {\\em perfectly} requires knowledge of the full distribution function (DF) of the system, or in other words, availability of a stationary solution of the collisionless Boltzmann equation.  However, such solutions are analytically known only for a very limited number of density distributions.  There is hence significant demand to construct equilibrium solutions numerically, not only for realizing N-body initial conditions but also in the context of modelling observational data sets.  In the latter case, finding such models is a main component of the reverse-engineering process aimed at constructing self-consistent three-dimensional systems that reproduce the observations. They can then be examined in great detail, allowing insights into properties that are not directly observable \\citep{Cretton1999, Bosch2012}.  To our knowledge, there are presently mainly five different methods in use for constructing such equilibrium models:  \\begin{enumerate} \\item DF-based: For certain mass distributions, the distribution function (DF) can be analytically calculated or accurately approximated.  Unfortunately, this ideal case is not generally available for arbitrary density distributions.  The main problem is that we do not know the analytical form of the third integral of motion. In some cases it may be reasonably approximated, but this leads at best to nearly self-consistent solutions \\citep{Kuijken1995, Widrow2005}.  Nevertheless, there are some useful classes of solutions known, for example for spherical galaxies \\citep{Osipkov1979, Jaffe1983, Merritt1985, Hernquist1990}. However, because many real systems are not particularly close to any of these parameterized classes of systems, the approach is rather restrictive in practice.  \\item Moment-based: Moments of the velocity distribution can be calculated or estimated with the hierarchy of Jeans equations. If one neglects higher order moments and assumes a functional form for the velocity distribution \\citep[often taken to be Gaussian, e.g.][]{Hernquist1993, Springel1999} that reproduces the estimated moments, one obtains an approximate distribution function.  This method is quite general and can be applied to all mass distributions. Since the true velocity distribution function is usually close to a triaxial Gaussian for much of the mass of a system, the method typically produces systems that are roughly in equilibrium. But the crux is that this equilibrium is by no means perfect, and that it is hard to overcome this limitation within this method. Especially difficult are the central regions of galaxies; when the constructed ICs are evolved in time, one here typically finds density ripples propagating through the system while it relaxes to a true equilibrium state. This can interfere with the interpretation of numerical experiments, especially when they require particularly quiet ICs.  \\item Orbit-based method: \\cite{Schwarzschild1979} introduced a radically different approach to solve the problem. He suggested to integrate a wide variety of orbits in a given potential, and then to distribute the mass of the system over this orbit library such that the time-averaged density of the system becomes as close as possible to the one corresponding to the potential.  Finding the weights of each of these orbits defines a linear optimization problem with positive coefficients, which can be solved iteratively. The resulting weights then effectively define the velocity distribution function.  A practical problem with this method is that the size of the orbit library is severely constrained by the available memory. Moreover, the method is ill-conditioned in its basic form, something that needs to be cured by adding ad-hoc assumptions such as smoothness constraints or maximum entropy measures for the weights. Also, the velocity distribution functions constructed with this method are typically very noisy and may feature large jumps. One needs to smooth them, but the required level of smoothing is hard to define. Many attempts have been made to overcome these difficulties \\citep[e.g.][]{Vandervoort1984, Jalali2011}.   \\item Made-to-measure: Attempts to improve on Schwarzschild's method have resulted in a new technique where the orbit integration process and the mass/weight redistribution are combined. This `made-to-measure' technique makes the storage of a full orbit library unnecessary and therefore removes the memory barrier. But it still requires a smoothing procedure for the velocity distributions \\citep{Syer1996, Dehnen2009}.  \\item Guided-relaxation: Another class of methods exploits the fact that any isolated system left to itself tends to an equilibrium state. Knowing the target mass distribution we may try to directly relax to it by steering a system appropriately in the process. This guiding can be done by introducing an additional force, e.g.~adiabatic drag on the vertical components of the particle velocities in order to squeeze the system \\citep{Holley-Bockelmann2001}, or we may restrict particle mobility such that the target density distribution is maintained and the systems evolves towards a self-consistent equilibrium state \\citep{Rodionov2009}. A disadvantage of this approach is that it involves one of the other methods to create an initial state for the further relaxation. Also, there is only limited control on the outcome, making it, e.g., difficult to construct systems with a prescribed velocity anisotropy.  \\end{enumerate}  The purpose of this article is to introduce a new, flexible approach for the construction of compound N-body models of axisymmetric galaxies in an essentially perfect equilibrium state. The method only requires the specification of the density profiles of the different components and a selection of the desired bulk properties of the velocity structure, such as the degree of rotational support or the ratio between radial and vertical velocity dispersion in the disk plane. Our code then constructs an N-body system that is in equilibrium and fulfills the imposed constraints on the velocity structure. Implicitly, it hence also provides a solution for the full 3D distribution function. This is achieved for essentially arbitrary axisymmetric density structure and by taking the mutual influence of different mass components (if present) fully into account.  We argue that the resulting flexibility and accuracy makes our approach an attractive alternative compared with other IC generation methods in the literature.  This paper is structured as follows. In Section~\\ref{sec:method}, we describe the basic methodology adopted in our method, which consists of an iterative procedure to adjust the velocities of an N-body realization of a galaxy model until the prescribed density structure is maintained self-consistently under time evolution, and the imposed velocity constraints are fulfilled. In Section~\\ref{sec:velconstraints}, we highlight how we specify velocity constraints for different structural choices. They take the form of second velocity moments which we determine through solutions of the Jeans equations. We then specify in Section~\\ref{sec:implementation} various implementation details of our numerical methods as realized in the {\\small GALIC} code introduced here. Section~\\ref{sec:galaxies} is concerned with a brief description of the specific density profile models currently implemented in this code; these are employed for a suite of tests presented in Section~\\ref{sec:testing}. Finally, we conclude with a summary of our findings in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we introduced a new iterative method for constructing equilibrium N-body galaxy models.  The principle advantage of our method is that it can produce N-body systems that are essentially in exact equilibrium for rather general density distributions, making the method ideal for studies of galaxy dynamics and numerical experiments with isolated or colliding galaxies. Compared to alternative schemes like the Schwarzschild method, our approach eliminates restrictions arising from a finite orbit library or from required regularization schemes.  Also, our method allows a natural inclusion of simulation aspects like the need for a gravitational softening.  The test results we have analyzed show a considerable improvement of the quality of the created initial conditions compared to existing codes such as the moment-based {\\small MAKENEWDISK}, a technique that has been used in numerous studies over recent years.  This is possible thanks to the absence of any assumptions in our approach with respect to the importance of higher-order moments of the velocity distribution function. The main disadvantage of our method lies in its higher computational cost compared to moment-based approaches. However, thanks to the scalable parallelization implemented of our code, this should not be a serious restriction in practice. For example it took about 5 hours on 96 AMD-6174 cores (2.2 GHz) to compute high-quality solutions for our most compicated models M1-4, while for the one component models it took only 1 hour. And since poor ICs may also impact any further scientific investigation, the additional CPU effort invested for better ICs should in many cases be well worth the effort.   Finally, we also note that numerous optimizations in our code could well be made to reduce its CPU time consumption. For example, the time to full convergence for a system with large $N$ may be reduced considerably by first treating a smaller subsample of the particles with correspondingly higher mass. Once this system has fully converged, one could then create the large target realization from it through bootstrap sampling, followed by briefly relaxing the big system to the final equilibrium.  In future work, it might be interesting to extend our approach to genuinely triaxial systems, which are of course considerably more challenging than the axisymmetric case considered here. One could even include additional phenomena such as figure rotation. An important challenge is here to suitably store the density response grid. Here, our approach, which only requires essentially one such response grid should be considerably less restrictive than Schwarzschild's method because the requirement to store a huge orbit library is avoided. In the meantime, we publicly release our {\\small GALIC} code (http://www.h-its.org/tap/galic), hoping that it proves useful for future N-body studies in galactic dynamics."
    },
    "1402/1402.1079_arXiv.txt": {
        "abstract": "{The {\\em Nuclear Spectroscopic Telescope Array (NuSTAR)} is the first focusing hard X-ray mission in orbit and operates in the 3--79 keV range. {\\em NuSTAR}'s sensitivity is roughly two orders of magnitude better than previous missions in this energy band thanks to its superb angular resolution. Since its launch in 2012 June, {\\em NuSTAR} has performed excellently and observed many interesting sources including four magnetars, two rotation-powered pulsars and the cataclysmic variable AE~Aquarii. {\\em NuSTAR} also discovered 3.76-s pulsations from the transient source SGR~J1745$-$29 recently found by {\\em Swift} very close to the Galactic Center, clearly identifying the source as a transient magnetar. For magnetar 1E~1841$-$045, we show that the spectrum is well fit by an absorbed blackbody plus broken power-law model with a hard power-law photon index of $\\sim$1.3. This is consistent with previous results by {\\em INTEGRAL} and {\\em RXTE}. We also find an interesting double-peaked pulse profile in the 25--35 keV band. For AE~Aquarii, we show that the spectrum can be described by a multi-temperature thermal model or a thermal plus non-thermal model; a multi-temperature thermal model without a non-thermal component cannot be ruled out. Furthermore, we do not see a spiky pulse profile in the hard X-ray band, as previously reported based on {\\em Suzaku} observations. For other magnetars and rotation-powered pulsars observed with {\\em NuSTAR}, data analysis results will be soon available.} ",
        "introduction": "A neutron star is the left-over dense and hot stellar core of a massive star which survived after the energetic explosion of a star. Neutron stars typically have strong surface magnetic fields ($10^9$--$10^{15}\\ \\rm G$) inferred from their spin properties. They are typically identified with pulsations in the radio, X-ray, or gamma-ray band which are caused by the rapid rotation of the star. There are many subclasses in the neutron star population such as rotation-powered pulsars (RPPs), rotating radio transients (RRATs), high-$B$ RPPs,  X-ray dim isolated neutron stars (XDINS), central compact objects (CCOs), and magnetars (see \\cite{k10} for a review).  Radiation properties of neutron stars are diverse. For example, RPPs are mostly observed in the radio band, but sometimes in the X-ray to gamma-ray band as well. Magnetars typically emit thermal photons (modified in the magnetosphere) in the soft X-ray band, but several show a striking turnover and rising power-law emission (in the $\\nu F_{\\nu}$ representation) in the hard X-ray band. X-ray properties of X-ray emitting neutron stars have been studied with the soft X-ray observatories which operate in the $\\sim$0.3--10 keV band (e.g., {\\em XMM-Newton} and {\\em Chandra}). However, the hard X-ray emission ($\\gapp$10 keV) has not been very well studied because until recently hard X-ray observatories could only detect a small handful of very bright sources.  While soft-band spectra have been well measured, they are strongly affected by the hydrogen absorption, and the current results based on the soft-band observations may be inaccurate. In particular, constraining a non-thermal spectral component is difficult because it can extend down to low energy and covaries with the hydrogen column density ($N_{\\rm H}$). Therefore, hard-band observations ($\\gapp$10 keV) which are almost free from the absorption can help us accurately determine the non-thermal spectrum.  \\label{sec:targets} \\begin{figure} \\includegraphics[width=80mm,height=60mm]{effectivearea.eps} \\caption{Effective area of {\\em NuSTAR} (red) together with that of {\\em XMM-Newton}/PN (blue) and {\\em Chandra}/ACIS (green).} \\label{fig:effarea} \\end{figure} \\medskip ",
        "conclusions": "\\label{sec:concl} Since its launch in 2012 June, {\\em NuSTAR} has performed excellently. As a part of {\\em NuSTAR}'s science program, we have observed carefully selected magnetars, RPPs and a white dwarf. Many interesting results are ready to be published for some sources, and detailed data analyses are on-going for others."
    },
    "1402/1402.1765_arXiv.txt": {
        "abstract": "We present time series photometric and spectroscopic data for the transient SN\\,2009ip from the start of its outburst in September 2012 until November 2013. This data was collected primarily with the new robotic capabilities of the Las Cumbres Observatory Global Telescope Network, a specialized facility for time domain astrophysics, and includes supporting high-resolution spectroscopy from the Southern Astrophysical Research Telescope, Kitt Peak National Observatory, and Gemini Observatory. Based on our nightly photometric monitoring, we interpret the strength and timing of fluctuations in the light curve as interactions between fast-moving ejecta and an inhomogeneous CSM produced by past eruptions of this massive luminous blue variable (LBV) star. Our time series of spectroscopy in 2012 reveals that, as the continuum and narrow H$\\alpha$ flux from CSM interactions declines, the broad component of H$\\alpha$ persists with SN-like velocities that are not typically seen in LBVs or SN Impostor events. At late times we find that SN\\,2009ip continues to decline slowly, at $\\lesssim0.01$ magnitudes per day, with small fluctuations in slope similar to Type IIn supernovae (SNe\\,IIn) or SN impostors, but no further LBV-like activity. The late time spectrum features broad calcium lines similar to both late time SNe and SN Impostors. In general, we find that the photometric and spectroscopic evolution of SN\\,2009ip is more similar to SNe\\,IIn than either continued eruptions of an LBV star or SN Impostors. In this context, we discuss the implications for episodic mass loss during the late stages of massive star evolution. ",
        "introduction": "\\label{sec:intro}  Type IIn supernovae (SN\\,IIn) are characterized by bright optical emission caused when explosion ejecta collides with nearby circumstellar material (CSM). This reveals that the progenitor star experienced a significant amount of mass loss prior to core collapse. The physical mechanism driving the mass loss, and the relative contributions to the CSM from winds and episodic eruptions, are not yet well constrained. In the final stages of massive star evolution, the star undergoes a rapid succession of burning phases, and exhibits variability during its mass loss episodes. The distribution of material in the CSM can be derived from spectroscopic and photometric fluctuations of SNe\\,IIn which, if the star was monitored for variability in previous years, allows for a connection between pre-explosion mass loss episodes and the material they produce. Luckily, the transient SN\\,2009ip affords us this rare opportunity.  SN\\,2009ip was first identified as a photometric transient near host galaxy NGC 7259 by the Chilean Automatic Supernova Search (CHASE; Maza et al. 2009), and given the designation `SN' without spectroscopic confirmation. Archival imaging of Hubble Space Telescope images of the progenitor star of SN\\,2009ip indicated it is very massive, $M=$ 50--80 $\\rm M_{\\odot}$ \\citep{Smith2010}. Analysis of archival data found a coincident variable source and suggested the transient may in fact be a luminous blue variable (LBV) or cataclysmic variable \\citep{Miller2009}. The LBV hypothesis was quickly confirmed with spectroscopy \\citep{Berger2009}, which showed SN\\,2009ip was similar to the group of ``SN imposters\" that look like Type IIn supernovae (e.g. Van Dyk et al. 2000). After peak brightness in 2009, the transient fell but then rose again in a manner similar to eta Carinae \\citep{Li2009}, and experienced a second outburst in mid-2010 \\citep{Drake2010}. Extensive analysis and discussion of the evidence for the 2009 and 2010 activity of SN\\,2009ip as an LBV eruption are provided by Smith et al. (2010) and Foley et al. (2011). In August 2012, SN\\,2009ip re-brightened to record luminosity \\citep{Drake2012}, maintaining an LBV-like spectrum \\citep{Foley2012}, but by mid-September 2012 the spectral lines had broadened and shifted, indicating material at velocities up to 13,000 $\\rm km\\ s^{-1}$, and showed P-Cygni profiles typical of core collapse supernova explosions \\citep{Smith2012}.  An analysis of spectroscopic and photometric data during the outburst of SN\\,2009ip from September to December of 2012 was presented in Mauerhan et al. (2013, hereafter JM13). They report Balmer lines with broad P-Cygni profiles and high absorption component velocities characteristic of core collapse supernovae, and completely unprecedented by any previous LBV eruption. JM13 observed SN\\,2009ip to fade, and then brighten to $M_R = -17.5$. The outbursts before and after this temporary fading are referred to as the 2012-A and 2012-B events (see Figure \\ref{fig:pastorello}). At the time of re-brightening the broad component of the Balmer lines disappeared, and only the narrow Lorentzian component remained. SN\\,2009ip then resembled a Type IIn supernovae, in which emission is dominated by the interaction of SN ejecta with the circumstellar medium (CSM). Based on their data, JM13 concluded that the 2012-A event of SN\\,2009ip was the core collapse of an LBV star during an outburst, and that the 2012-B event was caused by the fast-moving ejecta interacting with the CSM. This conclusion was supported by Prieto et al. (2013), who present a densely time-sampled optical light curve from the start of the 2012-B event and find that the rapid rise and bolometric luminosity are similar to other SNe\\,IIn.  \\begin{figure}[t] \\begin{centering} \\includegraphics[trim=0.7cm 0.0cm 0.0cm 0.0cm, clip, width=8.7cm]{fig_pastorello.eps} \\caption{Past eruptions of SN\\,2009ip as presented in Pastorello et al. (2013), and reproduced here for context. Lines connect the data points to guide the eye, but do not represent continuous photometric monitoring. \\label{fig:pastorello}} \\end{centering} \\end{figure}  The interpretation of SN\\,2009ip as a SN\\,IIn was challenged by Pastorello et al. (2013, hereafter AP13) who present three years of monitoring of SN\\,2009ip; we have included their photometry for context in Figure \\ref{fig:pastorello}. They find that spectra from September 24 2011 also exhibit material at velocity $\\sim13000$ $\\rm km\\ s^{-1}$, showing for the first time that non-terminal explosions can create such fast ejecta, and calling into question the supernova diagnosis for SN\\,2009ip. Instead, AP13 suggest the 2012 activity of SN\\,2009ip was caused by a pair-instability event during 2012-A, with the 2012-B event caused by ejecta material colliding with the CSM. The scenario of AP13 is fundamentally different from JM13 because the the star is not destroyed . This hypothesis is supported by the analysis of Fraser et al. (2013, hereafter MF13), who find a very low upper limit of $0.02 M_{\\odot}$ on the mass of any synthesized material (e.g. $\\rm ^{56}Ni$). A similar inference is made by Margutti et al. (2013, hereafter RM13), who analyze extensive multi-wavelength observations for SN\\,2009ip's 2012 activity, and conclude that the events were not caused by a core collapse but that the true physical mechanism behind the explosive ejection in the LBV's envelope could not, at that time, be fully distinguished. A full analysis including late-time observations of SN\\,2009ip by Smith et al. (2013, hereafter NS13) finds that the broad component of H$\\alpha$, first announced in Smith et al. (2012), persisted throughout the 2012-B event, indicating a significantly more energetic explosion has occurred than considered by RM13. They also find that the late-time spectrum of SN\\,2009ip no longer resembles an LBV. Combined, the observations of NS13 lead them firmly to the conclusion that a terminal supernova explosion has occurred.  This paper presents new, densely time-sampled photometry and spectroscopy for SN\\,2009ip from the The Las Cumbres Observatory Global Telescope Network (LCOGT.net), and other facilities at which we were observing classically at the time. LCOGT is a new system of robotic telescopes dedicated to time domain astrophysics. Telescopes of 1 and 2 meter diameter are distributed at five sites around the world, and observation requests are optimized and automatically executed by a real-time adaptive scheduler. Instrumentation includes identical imagers with a full suite of filters on all telescopes, and one robotic spectrograph (FLOYDS) on each of the two 2 meter Faulkes Telescopes. The distributed aspect of LCOGT make it flexible, responsive, and unhindered by weather or sunrise. New astrophysical transients are followed as soon as possible for classification, and can be reliably monitored on any desired time scale. The LCOGT Network, presented in depth by Brown et al. (2013), currently features the following facilities: McDonald Observatory, Texas USA (one 1m); Haleakala Observatory, Hawaii (one 2m+FLOYDS); Sutherland, South Africa (three 1m); Cerro Tololo, Chile (three 1m); and Siding Spring, Australia (two 1m and one 2m+FLOYDS).  The 2012--2013 activity of SN\\,2009ip is one of the first astrophysical transients followed with the LCOGT 1m network, and is a great example of the science potential from the combination of reactive follow-up and long term monitoring with a network that includes photometric and spectroscopic capabilities. The quality, quantity, and extent of these data allow us to provide new insight to the controversial physical interpretation of SN\\,2009ip, especially during it's extended slow decline. In \\S~\\ref{sec:obs} we present the observations; in \\S~\\ref{sec:LC} and \\ref{sec:spec} we interpret the photometry and spectroscopy respectively, and discuss the results in context with the ongoing debate about the true physical nature of SN\\,2009ip before concluding in \\S~\\ref{sec:con}. A flat cosmology with $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73$ is assumed throughout. ",
        "conclusions": "\\label{sec:con}  In this work we have used high cadence LCOGT data, with supporting observations from other facilities, to draw the following four main conclusions about SN\\,2009ip. (1) From our daily photometric monitoring, we observe photometric fluctuations from interactions between fast-moving ejecta and the CSM that reveal inhomogeneities in the CSM distribution. The light curve of the 2012-B event appears congruous with the timescales of past eruptions in 2011 and 2009. (2) The peak brightness and photometric evolution of SN\\,2009ip is more similar to SNe\\,IIn than LBVs or SN Impostors. Our long-term monitoring shows SN\\,2009ip continues to decline slowly with small fluctuations in the slope, which is shallower than if powered by the decay of Co$^{56}$. This is similar to both SNe\\,IIn and SN Impostors, but SN\\,2009ip has not exhibited any more LBV-like outbursts. (3) Our low and high resolution spectra during the 2012-B event exhibit evolving H$\\alpha$ emission and absorption components. As the continuum and narrow line flux from CSM interaction fades, the broad emission component persists. This is characteristic of SN-like expansion velocities, but not SN Impostors or LBVs. (4) The late-time spectrum of SN\\,2009ip exhibits forbidden calcium lines which are more similar to SNe\\,IIP, SNe\\,IIn, and SN Impostors than the LBV-phase spectrum of SN\\,2009ip. Asymmetry in the H$\\alpha$ line at late times hints at dust formation and/or asymmetry in the CSM, such as a disk.   SN\\,2009ip clearly fits into the class of SNe\\,IIn; if the progenitor has not suffered a terminal explosion, then perhaps not all SNe\\,IIn are actual supernovae. Multiple papers have remarked on the striking similarity -- right down to the timing and magnitude of a precursor outburst -- between SN\\,2009ip and the SN\\,IIn 2010mc (e.g. RM13, NS13). Far from a serendipitous match, Ofek et al. (2014) make a careful analysis of the Palomar Transient Factory data and find that such precursor events are common for SNe\\,IIn. SN\\,2009ip is unique in that it was observed first as a variable star undergoing eruptive mass loss for years before becoming a supernova. The resulting interactions of explosion ejecta with the eruption material reveal that the late stages of mass loss in massive stars can produce a CSM that has large inhomogeneities and a potentially asymmetric distribution.  SN\\,2009ip is a unique and interesting event for which multiple groups have assembled a high cadence photometric and spectroscopic coverage from a wide variety of telescopes. This publication contributes to the existing body of comprehensive data, with the main difference being that the bulk of this data has exceptionally regular cadence, and has been acquired with a single facility: the Las Cumbres Observatory Global Telescope, a network of robotic telescopes with imaging and spectroscopic capabilities. Only with the community's compilation of time series data is the true nature of SN\\,2009ip coming to be understood. High cadence multi-band and spectral coverage has traditionally been rare and challenging to accumulate, but the LCOGT Network can automatically monitor astronomical phenomena on almost any desired timescale; in the future, this can be done without human intervention. In particular, a distributed network like LCOGT mitigates losses due to weather and facilitates the acquisition of daily monitoring. LCOGT will continue to monitor SN\\,2009ip as it declines towards its fate."
    },
    "1402/1402.6619.txt": {
        "abstract": "We present the detection of hydrogen fluoride, HF, in two luminous nearby galaxies NGC~253 and NGC~4945 using the Heterodyne Instrument for the Far-Infrared (HIFI) on board the {\\it Herschel} Space Observatory. The HF line toward NGC~253 has a P-Cygni profile, while an asymmetric absorption profile is seen toward NGC~4945. The P-Cygni profile in NGC~253 suggests an outflow of molecular gas with a mass of {\\it M}(H$_2$)$_{out}$~$\\sim$ 1~$\\times$~10$^7$~M$_\\odot$ and an outflow rate as large as {\\it \\.{M}} $\\sim$~6.4~M$_\\odot$ yr$^{-1}$. In the case of NGC~4945, the axisymmetric velocity components in the HF line profile is compatible with the interpretation of a fast-rotating nuclear ring surrounding the nucleus and the presence of inflowing gas. The gas falls into the nucleus with an inflow rate of $\\le$~1.2~M$_\\odot$ yr$^{-1}$, inside a inner radius of $\\le$ 200 pc. The gas accretion rate to the central AGN is much smaller, suggesting that the inflow can be triggering a nuclear starburst. From these results, the HF $J = 1-0$ line is seen to provide an important probe of the kinematics of absorbing material along the sight-line to nearby galaxies with bright dust continuum and a promising new tracer of molecular gas in high-redshift galaxies.   %The HF abundances with respect of H$_2$ of a few 10$^{-9}$, calculated toward both sources, are one order of magnitude smaller than those observed toward diffuse gas in the Milky Way, but also one order of magnitude larger than those observed in galactic massive star formation regions. The high abundance observed in dense and dynamically active regions can be explain by HF thermal desorption from the ice grains. %{\\it M}(H$_2$)$_{out}$ $=$ 2.5 $\\times$ 10$^7$ M$_\\odot$; the molecular mass loss rate \\.{M} $=$ 13.3 M$_\\odot$ yr$^{-1}$; maximum outflow velocity of {\\it V}$\\rm_{out}$ = 157 km~s$^{-1}$  %are due to $\\sim$ 160 km $s^{-1}$ %The spectrum obtained shows a complex pattern of %tion, with numerous features covering a wide range of local standard of rest velocities (?130 to 100 km ?1). An analysis of this absorption yields HF abundances relative to H2 of ?1.3 ? 10?8, in most velocity intervals. %The Heterodyne Instrument for the Far-Infrared (HIFI) on board the Herschel Space Observatory (HSO) has provided a unique opportunity to study hydride molecules with large rotational constant in the interstellar medium (ISM) within the Milky Way and nearby galaxies. Within hydride molecules, hydrogen halides are particularly interesting because of its unique thermochemistry. One key result from Herschel/HIFI instrument observations is the first detection of the fundamental $J = 1-0$ rotational transition of hydrogen fluoride (HF) and the discovery of its ubiquitous nature within the ISM of the Milky Way. Because of the thermochemical properties of fluorine, HF becomes the dominant reservoir of gas-phase fluorine in regions where the gas is primarily molecular and thus HF can be an excellent surrogate for H2.  %These and other interesting results on hydrogen halides from Herschel observations will be presented, as well as the results from follow-up studies from ground-base facilities on hydride molecules toward luminous lensed high-redshifted galaxies with strong submillimeter continuum ",
        "introduction": "%\\textcolor{red}{Tom: Do you have a case for CO (or H I) where you can show what an infall with an AGN looks like, compared to a starburst?}  %\\textcolor{red}{Darek: win you compute Gamma c or l, do you include the emission component as additional background?}  %Hydride molecules, compounds form with one or more hydrogen and any other element from the periodic table, are particularly important in astrochemistry why?: since they are  {\\it Herschel} Space Observatory \\citep{Pil10}, and in particular its Heterodyne Instrument for Far-Infrared (HIFI) \\citep{deG10} has facilitated observations at high-spectral resolution of interstellar hydride molecules - compounds with one or more hydrogen and one heavy element atom -, in the local universe (z $\\sim$ 0). Most of the hydride molecule observations are inaccessible from ground-based telescopes because their high frequency rotational transitions are blocked by the large opacity of Earth's atmosphere. Two key results from {\\it Herschel}/HIFI were the first detection of the \\emph{fundamental} $J=1-0$ rotational transition of hydrogen fluoride (HF) at 1.232 THz (243 $\\mu$m) and the discovery of HF's ubiquitous nature in the Milky Way. HF has not only been observed in almost every bright continuum source in the Galactic plane \\citep{Phi10,Son10,Neu10,Mon11a}, but also in some nearby ultra luminous galaxies \\citep{van10,Ran11}, establishing its importance outside the Milky Way as well. %HF is form after the exothermal reaction of fluorine and H$_2$ %Fluorine can be found mainly in its neutral form in the diffuse ISM %HF chemistry: Despite Fluorine's (F) relative low abundance in the interstellar medium (ISM) (about four times lower than carbon), F plays an important role in the interstellar chemistry due to the unique thermochemistry of the reaction between F and molecular hydrogen (H$_2$). F is the only atom that reacts exothermically with H$_2$, to form the compound HF. Once formed, HF becomes the main reservoir of fluorine in the ISM, with a strong bond only destroyed by reactions with low abundance ions H$^{+}_{3}$, C$^{+}$ and He$^{+}$, or photodissociation, with an estimated photodissociation rate of 1.17~10$^{-10}$~s$^{-1}$. This unusual stability allows the build up of large amounts of HF in the ISM, which has now been confirmed by {\\it Herschel}.  %\\textcolor{red}{Hydrogen fluoride (HF) is expected to be the main reservoir of fluorine (F) in the interstellar medium (ISM) because of its unique thermochemistry. Fluorine, predominantly neutral in the diffuse ISM, reacts exothermically with H$_{2}$, the dominant constituent of molecular clouds, forming HF. Once the H$_{2}$ abundance becomes appreciable, HF becomes the dominant reservoir of fluorine and is destroyed only slowly by photodissociation, with an estimated photodissociation rate of 1.17~10$^{-10}$~s$^{-1}$ for the mean {\\bf interstellar} radiation field, and by reactions with He$^{+}$, H$^{+}_{3}$, and C$^{+}$. }   Theoretical models from \\cite{Neu05} originally predicted that the HF abundance in diffuse clouds of small extinction ({\\it A$_V$} $<$ 0.5), could be larger than that of the carbon monoxide (CO) - the main tracer of molecular hydrogen in the submillimeter range- despite the lower abundance of gas-phase fluorine compared to that of carbon. Furthermore, the models also showed that HF may be a more reliable tracer of H$_2$ than CO, since it presents a constant HF/H$_2$ ratio (see Figure 6 in Neufeld et al. 2005) with respect to the total visual extinction through the cloud, while the CO/H$_2$ ratio decreases in clouds of small {\\it A$_V$}. \\\\ Thus, the \\cite{Neu05}  model predicted that the ground state rotational transition line of HF $J=1-0$ would yield an extremely sensitive probe of the diffuse molecular gas along the line-of-sight to background far-infrared continuum sources and a potential valuable surrogate for molecular hydrogen. %HF thus becomes the main reservoir of fluorine in the interstellar medium. %\\textcolor{red}{To date, such material has been studied in the submillimeter range primarily by means of carbon monoxide (CO) rotational emission lines. However, \\cite{Neu05} show in their models that, in diffuse clouds of small extinction, the predicted HF abundance can actually exceed that of CO, even though the gas-phase fluorine abundance is 4 orders of magnitude smaller than that of carbon. HF is a more reliable tracer than CO because the HF/H$_{2}$ ratio is constant \\citep[see Figure~6 in][]{Neu05}, whereas CO/H$_{2}$ drops in clouds of small {\\it A$_{\\rm v}$}.} The high Einstein coefficient for spontaneous emission coefficient of the HF $J=1-0$ line, {\\it A$_{ \\rm 10}$}~=~2.42~10$^{-2}$~s$^{-1}$, results in simple excitation, with very high critical density ($\\sim$ 10$^9$ cm$^{-3}$), at which collisional processes become important for the excitation. Thus, under typical conditions characteristic of the diffuse or even dense ISM, the HF molecules are mainly in the ground rotational state. As a result of its large {\\it A$_{10}$} coefficient the HF $J=1-0$ line has also been observed primarily in absorption. Only an extremely dense region or a strong radiation field could generate enough excitation to yield an HF feature with a positive frequency-integrated flux. %\\textcolor{red}{ Thus, essentially all HF molecules are in the ground rotational state under typical conditions characteristic of the diffuse or even dense ISM.} %{\\bf Observations in the local universe of the fundamental HF $J=1-0$ rotational transition at 1.232 THz (243 $\\mu$m) has been achievable for %the first time with the use of the Heterodyne Instrument for Far-Infrared (HIFI) \\citep{deG10} on board the {\\it Herschel} Space Observatory \\citep{Pil10}.} %\\textcolor{blue}{The Heterodyne Instrument for Far-Infrared (HIFI) \\citep{deG10} on board the {\\it Herschel} Space Observatory \\citep{Pil10} has allowed for the first time the detection of the \\emph{fundamental} $J=1-0$ rotational transition of HF at 1.232 THz (243 $\\mu$m) in the local universe. HF has proven to be a ubiquitous tracer of molecular gas in the ISM of the Milky Way, being detected in environments as diverse as Orion~KL, OMC-1 \\citep{Phi10}, as well as diffuse clouds on the line-of-sight toward W49N, W51 \\citep{Son10}, W31C \\citep{Neu10}, Sgr B2 (M) \\citep{Mon11a} and in some nearby active galaxies, such as, Mrk 231 \\citep{van10} and Arp 220 \\citep{Ran11}. } %\\textcolor{red}{The HF $J=1-0$ transition is generally observed in absorption, as expected, due to its very large {\\it A} coefficient. Only an extremely dense region or a strong radiation field could generate enough excitation to yield an HF feature with a positive frequency-integrated flux. } \\cite{van12} present HF observations toward the Orion Bar,  where the line appears in emission and demonstrate that the excitation of HF is dominated by collisions with electrons, and argue that similar conditions can be found toward active galaxies such as Mrk 231 \\citep{van10}, where the radiation field is strong and the HF line also appears in emission.      The HIFI observations support the theoretical predictions that HF will be the dominant reservoir of interstellar fluorine \\emph{under a wide range of conditions}. The most interesting aspect is that HF will likely be an important tracer of molecular gas in high-redshift galaxies, as seen by recent observations with the Caltech Submillimeter Observatory (CSO) that revealed the highest-redshift detection of interstellar HF to date toward the luminous lensed Cloverleaf galaxy at {\\it z} = 2.558 \\citep{Mon11b}. High-redshift HF observations will be done routinely with the large collecting area of the Atacama Large Millimeter/submillimeter Array (ALMA). However, to qualitatively analyze the distant galaxy measurements, a good understanding of HF in the galaxies in the local universe is needed, which to date, is only achievable with {\\it Herschel}. In this paper, we study the HF content of two nearby galaxies NGC~253 and NGC~4945, prototypical examples of a nearby starburst and a composite active galactic nucleus (AGN) -- starburst nuclei, respectively. NGC~253 is an edge-on ({\\it i} = 78$^{\\circ}$) nearby barred galaxy associated with the Sculptor group at a distance of {\\it D}~=~3.5~Mpc, having an IR luminosity of {\\it L}$_{FIR}$ = 1.7 $\\times$ 10$^{10}${\\it L}$_{\\odot}$ \\citep{Rad01} originated in intense massive star formation regions within its central few hundred parsecs \\citep{Str04}. It is also suggested that NGC~253 contains a weak AGN in conjunction with the strong starburst \\citep[e.g.][]{Mul10}. NGC~4945 is also a nearly edge-on ({\\it i} = 78$^{\\circ}$) local galaxy, part of the Centaurus group at a distance of $\\sim$~3.8 Mpc \\citep{Kar07}, with a total infrared luminosity of {\\it L}$_{FIR}$ = 2.4 $\\times$ 10$^{10}${\\it L}$_{\\odot}$ \\citep{Bro88}. %NGC~4945 host both an AGN and a starburst. It is believed to be a post starburst phase (koornneef 1993) or in a transition from a starburst to a Seyfert galaxy (Moorwood \\& Oliva 1994). \\textcolor{red}{more...}  In Sections 2 and 3 the observations and results are described. In Section 3.1 we obtain the HF column densities for each galaxy. Outflow and inflow gas properties and their nature are discussed in Section 3.2 and 3.3, with conclusions in Section 4.   % where the HF and H$_2$ column densities would track each other closely with the ratio of {\\it N}(HF)/{\\it N}(H$_2$) equal to the gas phase elemental abundance of F relative to H nuclei \\citep[equal to 3.6 $\\times $10$^{-8}$ in diffuse atomic clouds in the solar neighborhood][]{Sno07}. %Herschel observations of HF $J = 1-0$ toward diverse Galactic environments  the derived abundances, with an average ratio of {\\it N}(HF)/{\\it N}(H$_2$)=(1.44~$\\pm$~0.35)~$\\times$~10$^{-8}$ \\citep{Son10,Mon11b}, are a factor of two below the chemical model predictions. %Chemical analysis of fluorine-containing molecules in the ISM \\citep[e.g.][]{Neu05} have predicted that over a wide range of conditionsHerschel observations of HF $J = 1-0$ toward diverse Galactic environments proved that HF is an excellent probe of H$_2$ in diffuse clouds  %To quantitatively analyze the distant galaxy measurements, good understanding of hydride molecules in the local universe is also needed. This is the main goal of our Open Time 2 (OT2) proposal for Herschel/HIFI, awarded with 25.2 hours, to study the hydrogen fluoride in a sample of 16 nearby LIRGs. The sources we proposed will be used as templates in the interpretation of the future observations of HF in high-redshift galaxies.  %Add something on NGC~253 and NGC~4945.  %NGC~253, a nearly edge-on spiral ({\\it i} = 78$^{\\circ}$) %from Aalto: %NGC~253 is also a nearby barred galaxy located in the Sculptor group with a compact nuclear starburst and a IR luminosity that appears to originate in regions of intense %massive star formation within its central few hundred parsecs (Strickland et al. 2004). From their 2 mm spectral scan Mart\u00ab\u00f5n et al. (2006) suggest that the chemistry of NGC~253 shows strong similarities to that of the Galactic Center molecular region, which is thought to be dominated by low-velocity shocks. High resolution SiO observations show bright emission resulting from large scale shocks as well as gas entrained in a nuclear out\u00dfow (Garc\u00ab\u00f5a-Burillo et al. 2000). Note also that it is suggested that NGC~253 is a galaxy in which a strong starburst and a weak AGN coexist (e.g. Weaver et al. 2002; M\u00aculler-S\u00abanchez et al. 2010). %end of Aalto.  %Mauersberger: %The Sculptor galaxy NGC~253, a nearly edge-on spiral ({\\it i} = 78$^{\\circ}$) of type SAB(s)c perhaps the best example for a nearby nuclear starburst. %end Maue %In NGC~253, a bright edge-on Sc galaxy associated with the Sculptor group. ",
        "conclusions": "We present the detection of the HF $J=1-0$ transition in two nearby galaxies, NGC~253 and NGC~4945. The HF spectrum toward NGC~253 shows a P--Cygni profile, with the absorption blue-shifted with respect to the systemic velocity, as expected for an outflow system. Estimated outflow parameters derived from the HF spectrum are mass of the outflowing gas {\\it M}(H$_2$)$_{out}$ $\\sim$ 1 $\\times$ 10$^7$ M$_\\odot$; the molecular mass loss rate {\\it \\.{M}} $\\le$ 6.4 M$_\\odot$ yr$^{-1}$; and a maximum unprojected outflow velocity of {\\it $\\nu$}$\\rm_{out}$ $\\sim$ 194~km~s$^{-1}$. The relatively small outflow velocity and the SFR comparable to the outflow rate, argue for a starburst nature of the outflow, however other driving mechanisms, such as radiation pressure, cannot be ruled out. While the HF spectrum shows a clear evidence of a molecular outflow, the conclusions about its nature need further confirmation with spatially-resolved observations or a larger statistical sample. The HF line toward NGC~4945 shows an asymmetric absorption feature, where we can distinguish two components at almost axi-symmetric velocities with respect to the systemic velocity ($\\nu$$_ {\\rm \\scriptscriptstyle SYS}$ $\\pm$ 50 km~s$^{-1}$). The axisymmetry velocity profile is compatible with the ring and inflowing gas scenarios, where the redshifted component can be interpreted as a signature of gas falling onto the nucleus. The relatively high inflow rate obtained from the absorption spectrum, indicates gas triggering a nuclear starburst. However, higher angular resolution are needed to reach a firmer conclusion about the nuclear kinematics of NGC~4945. %While the HF spectrum shows a clear evidence of non circular motion indicating molecular outflow and inflow, the conclusions about its nature need further confirmation with spatially resolved observations or a larger statistical sample.  %The derived HF abundances with respect to H$_2$, are 1.13~$\\times$~10$^{-9}$ and 2.44~$\\times$~10$^{-9}$ toward NGC~253 and NGC~4945, respectively. These values are one order of magnitude lower than those observed toward diffuse galactic sources. Based on the average fluorine (F) gas-phase interstellar abundance in diffuse atomic gas clouds, {\\it N}$_F$/{\\it N}$_H$ = 1.8 $\\times$ 10$^{-8}$ \\citep{Sno07}, our observations suggest a F depletion factors of $\\sim$~25 and 13, respectively. Lower values of HF abundances by about two orders of magnitude have been obtained in high-mass star-forming regions, with much higher densities compared to diffuse clouds \\citep[e.g.][observed HF abundances as low as 5~$\\times$~10$^{10}$ in the denser parts of the massive star formation region of NGC 6334 I]{Emp12}. The low abundance of HF under these conditions is most likely caused by freeze--out of HF molecules onto dust grains. This argument is derived from the fact that almost all fluorine should be bound in HF due to its high proton affinity, and that the desorption energy of HF is assumed to be quite large due to its polar nature. In dense (10$^5$~-~10$^6$~cm$^{-3}$) and warm (100~-~150 K), but dynamically active, regions such as outflows or inflows, we find HF abundances higher than in dense star--forming regions, but in our case, still lower than values found in diffuse clouds. In the picture of HF freeze--out at high densities, this can be explained by thermal desorption of HF from the ice grains, similar to the desorption of water in outflows \\citep{Fra08,Kri10}.  From these results, the HF $J=1-0$ line promises to provide a valuable probe of the kinematics of absorbing material along the sight-line to bright extragalactic continuum sources. HF $J=1-0$ line observations towards a larger sample of nearby (U)LIRGs, including radiative transfer models will be presented in a future publication. Studies with larger statistical samples will be able to establish a correlation between the HF abundance and the physical conditions, especially density and temperature. The characterization of HF line observations towards nearby galaxies will be an important tool to understand future observations of HF at high redshift, where the HF line could be a relevant probe of molecular gas in dense but especially in lower density regions of those high redshifted sources where the redshifted H \\textsc{i} frequency falls outside the frequency allocation windows for radio astronomy.   %to determine the nature of the HF depletion.  %In the high-mass star-forming regions, at much higher densities compared with diffuse clouds, the HF abundance seems to drop by about two orders of magnitude toward several Galactic sources [REF, REF]. %The HF column densities calculated from the absorption profile  %These agrees also with chemical studies of fluorine-containing molecules in the ISM \\citep[e.g.][]{Neu05} that predicted that over a wide range of conditions the HF and H$_2$ column densities would track each other closely with the ratio of {\\it N}(HF)/{\\it N}(H$_2$) equal to the gas phase elemental abundance of F relative to H nuclei \\citep[equal to 3.6 $\\times $10$^{-8}$ in diffuse atomic clouds in the solar neighborhood][]{Sno07}.  %\\end{enumerate}  %Chemical analysis of fluorine-containing molecules in the ISM \\citep[e.g.][]{Neu05} have predicted that over a wide range of conditions the HF and H$_2$ column densities would track each other closely with the ratio of {\\it N}(HF)/{\\it N}(H$_2$) equal to the gas phase elemental abundance of F relative to H nuclei \\citep[equal to 3.6 $\\times $10$^{-8}$ in diffuse atomic clouds in the solar neighborhood][]{Sno07}. Herschel observations of HF $J = 1-0$ toward diverse Galactic environments proved that HF is an excellent probe of H$_2$ in diffuse clouds, although the derived abundances, with an average ratio of {\\it N}(HF)/{\\it N}(H$_2$)=(1.44~$\\pm$~0.35)~$\\times$~10$^{-8}$ \\citep{Son10,Mon11b}, are a factor of two below the chemical model predictions. Lower values of HF abundances have been obtained toward denser regions due to the freeze out of HF onto dust grains \\citep[e.g.][observed HF abundances as low as 5~$\\times$~10$^{10}$ in the denser parts of the massive star formation region NGC6334 I]{Emp12}. The detection of the HF absorption line toward nearby galaxies source shows that the HF line is also an extremely sensitive tool for tracing molecular material toward z~$\\ge$~0 objects and their gas kinematics. %The Herschel/HIFI observational results of HF $J=1-0$ toward the nearby galaxy NGC~253 shows that the HF line is a good tracer of molecular outflows. The derived abundance of HF with respect to H$_2$ is 5.6~$\\times$~10$^{-10}$. Based on the average gas-phase interstellar abundance in diffuse atomic gas clouds mention above, {\\it N}$_F$/{\\it N}$_H$ = 1.8 $\\times$ 10$^{-8}$ \\citep{Sno07}, our observations suggest a F depletion factor of $\\sim$~10     %\\clearpage"
    },
    "1402/1402.0625_arXiv.txt": {
        "abstract": "Both Pluto and its satellite Charon have rotation rates synchronous with their orbital mean motion.  This is the theoretical end point of tidal evolution where transfer of angular momentum has ceased.  Here we follow Pluto's tidal evolution from an initial state having the current total angular momentum of the system but with Charon in an eccentric orbit with semimajor axis $a \\approx 4R_{\\s P}$ (where $R_{\\s P}$ is the radius of Pluto), consistent with its impact origin. Two tidal models are used, where the tidal dissipation function $Q\\propto$ 1/frequency and $Q=$ constant, where details of the evolution are strongly model dependent. The inclusion of the gravitational harmonic coefficient $C_{22}$ of both bodies in the analysis allows smooth, self consistent evolution to the dual synchronous state, whereas its omission frustrates successful evolution in some cases.  The zonal harmonic $J_2$ can also be included, but does not cause a significant effect on the overall evolution.  The ratio of dissipation in Charon to that in Pluto controls the behavior of the orbital eccentricity, where a judicious choice leads to a nearly constant eccentricity until the final approach to dual synchronous rotation. The tidal models are complete in the sense that every nuance of tidal evolution is realized while conserving total angular momentum --- including temporary capture into spin-orbit resonances as Charon's spin decreases and damped librations about the same. ",
        "introduction": "Pluto has five known satellites: Charon, Nix, Hydra, Keberos, and Styx, with the latter four much smaller than Charon. Listed in Table~\\ref{para} are the physical and orbital parameters of Pluto-Charon from \\cite{Buie12}, unless otherwise specified. The Charon-Pluto mass ratio ($q=0.1165$) is large when compared with others in the Solar System ($1/81$ for Moon-Earth and $< 1/4000$ for the other satellites and their planets). The barycenter of the Pluto-Charon system lies outside the surface of Pluto. Hence, some astronomers regard the pair as a binary system \\citep{Stern92}. The total angular momentum $L$ of the Pluto-Charon system is so large that the combined pair would be rotationally unstable \\citep{Mignard81a,Lin81}.  The Pluto-Charon system is currently in a dual synchronous state \\citep{Buie97,Buie10}, which is the endpoint of tidal evolution. As such the expected zero orbital eccentricity has been recently verified (with a 1-$\\sigma$ upper limit of $7.5 \\times 10^{-5}$), after taking into account the effects of surface albedo variations on Pluto (\\citealt{Buie12}; see Table~\\ref{para}).   \\begin{deluxetable}{lcc} \\tablecolumns{3} \\tablewidth{0pt} \\tablecaption{Physical and orbital parameters of the Pluto-Charon system \\label{para}} \\tablehead{ \\colhead{Parameter} & \\colhead{Pluto} & \\colhead{Charon} } \\startdata $GM$ (km$^3$\\,s$^{-2}$)\\tablenotemark{a} & 870.3(3.7) & 101.4(2.8) \\\\ Radius $R$ (km)\\tablenotemark{b} & 1153(10) & 606.0(1.5) \\\\ Orbital period $P$ (days) & \\multicolumn{2}{c}{6.3872273(3)} \\\\ Semimajor axis $a$ (km) & \\multicolumn{2}{c}{19573(2)} \\\\ Eccentricity $e$ &\\multicolumn{2}{c}{0} \\\\ Inclination $i$ ($^{\\circ}$)&\\multicolumn{2}{c}{96.218(8)} \\\\ Long.\\ ascending node $\\Omega$ ($^{\\circ}$) & \\multicolumn{2}{c}{223.0232(69)} \\\\ \\enddata \\tablenotetext{a}{Adopted from \\cite{Tholen08}, where $G$ is the Newtonian gravitational constant and $M$ is the mass.} \\tablenotetext{b}{Adopted from \\citet{Buie06} for Pluto and \\citet{Person06} for Charon.} \\tablecomments{ The orbital elements are Pluto-centric with respect to the mean equator and equinox of J2000 at the epoch JD 2452600.5. Numbers in parentheses are 1-$\\sigma$ errors in the least significant digits.} \\label{parameters} \\end{deluxetable}  As Pluto-Charon is similar to Earth-Moon, the feasible origin of this system may be chosen from the proposed schemes for the origin of the Earth-Moon system.  A giant impact of a Mars-sized body is thought to be the only viable origin of the Moon  (e.g., \\citealt{Cameron76,Boss86,Canup04}) to account for the large angular momentum of the system. \\cite{McKinnon84} proposed a similar origin for Charon.  If Charon accumulated from a debris disk resulting from such an impact, the initial eccentricity of Charon's orbit would be near zero. \\citet[][hereafter DPH97]{Dobrovolskis97} were thereby motivated to determine the tidal evolution of Charon in a circular orbit to the current dual synchronous state in a time short compared to the age of the Solar System (see also \\citealt{Farinella79}) as the only possible outcome of the dissipative process. In a circular orbit, Charon would reach synchronous rotation very quickly (e.g., DPH97), and this has generally been assumed (e.g., \\citealt{Peale99}).  However, smoothed particle hydrodynamic (SPH) simulations by \\cite{Canup05} showed that the results of a nearly intact capture in a glancing encounter surround the $(q,L)$ region of the system much more completely than those of disk-forming impacts. Therefore, capture where Charon comes off nearly intact after a glancing impact is favored and non-zero eccentricity would be more probable.  We are not aware of any previous attempts to examine the tidal evolution of Charon's orbit incorporating finite eccentricity. As we shall see, Charon in an initially eccentric orbit avoids the almost immediate synchronous rotation heretofore assumed, and the varied and interesting evolutionary sequences that were suppressed in the circular orbit evolution are revealed. Depending on the ratios of rigidity $\\mu$ and tidal dissipation function $Q$ between Pluto and Charon, the eccentricity of Charon's orbit may either grow or decay during most of the evolution \\citep{Ward06}. Permanent quadrupole moments of the bodies may also lead to spin-orbit resonance, and such resonances can have a significant effect on the orbital evolution.  In the following we tidally evolve the Pluto-Charon system with two tidal models distinguished by the dependence of the dissipation function $Q$ on frequency $f$: $Q\\propto 1/f$ and $Q=$ constant. The tidal model developed in Section~\\ref{Deltatmodel} has the tidal distortion of a body responding to the perturbing body a short time $\\Delta t$ in the past. Constant $\\Delta t$ leads to  $Q\\propto 1/f$, so we call the $Q\\propto 1/f$ model the constant $\\Delta t$ model. In Section~\\ref{Qmodel} we develop the equations of evolution for the constant $Q$ model. Although neither of these frequency dependences represent the behavior of real solid materials (e.g., \\citealt{Castillo11}) and although the evolutionary tracks are model dependent, most if not all of the possible routes from probable initial configurations to the current equilibrium state are demonstrated. In Section~\\ref{J2C22} we develop the contributions of rotational flattening $J_2$ and permanent quadrupole moment $C_{22}$ to the equations of motion. We describe the adopted system parameters and initial conditions in Section~\\ref{ICs} and the numerical methods in Section~\\ref{numerics}. The results from both the constant $\\Delta t$ and constant $Q$ models with zero $J_{2{\\s P}}$ for Pluto and zero $C_{22}$ for both bodies are shown in Section~\\ref{results1}, and the effects of non-zero $J_{2P}$ and $C_{22}$ in Section~\\ref{results2}, respectively. The results are discussed in Section~\\ref{discussion}, and the conclusions are summarized in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have investigated the tidal evolution of Pluto-Charon on an eccentric orbit under two different tidal models: constant $\\Delta t$ and constant $Q$. Our calculations show the complete tidal evolution of a system of two solid bodies of comparable size, where the spin angular momentum of the two bodies is initially comparable to the orbital angular momentum. The deviation from axial symmetry has been included in tidal evolution, and the back reaction on the orbit must be accounted for to conserve angular momentum. Capture into spin-orbit resonances can profoundly affect the tidal evolution of the system.  Motivated by binary asteroids (including those with comparable masses in dual synchronous state, like (69230) Hermes and (90) Antiope), \\cite{Taylor10,Taylor11} have studied the tidal evolution of two bodies. They included higher order terms in the tidal potential, but considered only circular orbit and zero $C_{22}$. \\cite{Rodriguez12} have studied the tidal evolution of super-Earths close to their host stars. They considered non-zero $C_{22}$ and eccentric orbit, and also found that capture into spin-orbit resonances can significantly affect the tidal evolution of eccentricity. However, they included only the tides raised by the star on the planet. Both of these studies used the constant $\\Delta t$ model only.  The equations used in our study are derived from a variety of existing sources, and we provide a more comprehensive listing of the coefficients in the evolution equations of the constant $Q$ model, which depend discontinuously on the spin rate. In both models, we find the value and range of relative rates of tidal dissipation in Charon to that in Pluto that would result in roughly constant eccentricity during most of the evolution. In the constant $\\Delta t$ model, the results are valid for arbitrary eccentricity, which is not true for constant $Q$ (where the results are qualitatively inaccurate for $e \\gtrsim 0.36$ due to the necessary truncations in the evolution equations). However, the constant $Q$ model requires a more reasonable relative rate of dissipation between Pluto and Charon ($A_Q \\sim 1$).  It was assumed in previous studies (e.g., DPH97) that Charon would achieve synchronous rotation quickly after its formation. We show that this is not the case for Charon on an eccentric orbit. The asymptotic spin depends on both the eccentricity and the assumed tidal model. While the inferred large oblateness of Pluto gives no significant change to the evolution, it is found that the capture into spin-orbit resonance of Charon for non-zero values of $C_{22}$ can change the relative dissipation rate that keeps the eccentricity more or less constant during most of the evolution. In some cases (e.g., if $A_{\\Delta t} \\sim 1$), spin-orbit resonance can allow smooth evolution to the final state of dual synchronous rotation, whereas very large eccentricity and semimajor axis would otherwise occur (which could lead to instability). The conditions of capture into and escape from the 3:2 spin-orbit resonance as a function of the orbital eccentricity agree with the existing results in the literature."
    },
    "1402/1402.4026.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We present a study of the environment of the Swift long gamma-ray burst GRB\\,120327A at $z\\approx 2.8$ through optical spectroscopy of its afterglow.} % methods heading (mandatory) {We analyzed medium-resolution, multi-epoch spectroscopic observations ($R \\sim 7000\\mbox{--}12000$, corresponding to $\\sim 15\\mbox{--} 23$ km~s$^{-1}$, $\\mbox{S/N} = 15 \\mbox{--} 30$ and wavelength range 3000--25000~\\AA) of the optical afterglow of GRB\\,120327A, taken with X-shooter at the VLT $2.13$ and $27.65$ hr after the GRB trigger.} % results heading (mandatory) {The first epoch spectrum shows that the ISM in the GRB host galaxy at $z = 2.8145$ is extremely rich in absorption features, with three components contributing to the line profiles. The hydrogen column density associated with GRB\\,120327A has $\\log N_{\\rm H}/{\\rm cm}^{-2} = 22.01 \\pm 0.09$, and the metallicity of the host galaxy is in the range $\\mbox{[X/H]} = -1.3$ to $-1.1$. In addition to the ground state lines, we detect absorption features associated with excited states of \\ion{C}{ii}, \\ion{O}{i}, \\ion{Si}{ii}, \\ion{Fe}{ii}, and \\ion{Ni}{ii}, which we used to derive information on the distance between the host absorbing gas and the site of the GRB explosion. The variability of the {\\ion{Fe}{ii}~$\\lambda$2396} excited line between the two epochs proves that these features are excited by the GRB UV flux. Moreover, the distance of component I is found to be $d_{\\rm I}=200^{+100}_{-60}$~pc, while component II is located closer to the GRB, at $d_{\\rm II}=100^{+40}_{-30}$~pc. These values are among the lowest found in GRBs. Component III does not show excited transitions, so it should be located farther away from the GRB.  The presence of H$_2$ molecules is firmly established, with a molecular fraction $f$ in the range $f=4\\times 10^{-7} \\mbox{--} 10^{-4}$. This particularly low value can be attributed to the small dust content.  This represents the third positive detection of molecules in a GRB environment. %The GRB\\,120327A host galaxy %    appears dusty ([Zn/Fe]$=0.55\\pm0.13$), but not as in the dustiest %    GRB hosts, and chemical enrichment models suggest that it is an %    elliptical with a mass of $\\sim 10^{11} M_{\\odot}$.  %A comparison with %    galactic chemical evolution models has suggested that the host of %    GRB\\,090926A is likely to be a dwarf-irregular galaxy. No emission %    lines were detected, but a H$\\alpha$ flux in emission of %    $9\\times10^{-18}$ erg~s$^{-1}$~cm$^{-2}$ (i.e., a star-formation %    rate of $2~M_\\odot$yr$^{-1}$), which is typical of many GRB hosts, %    would have been detected in our spectra, and thus emission lines %    are well within the reach of X-shooter. We put an upper limit to %    the \\ion{H} molecular fraction of the host galaxy ISM, which is $f %    < 7\\times10^{-7}$.  The continuum has been fitted assuming a %    power-law spectrum, with a spectral index of $\\beta = %    0.89^{+0.02}_{-0.02}$. The best fit does essentially not require %    any intrinsic extinction because $E_{B-V} < 0.01$ mag adopting a %    SMC extinction curve. Finally, the line of sight of GRB\\,090926A %    presents four weak intervening absorption systems in the redshift %    range $ 1.24 < z < 1.95$. %Three %of them feature the {\\ion{C}{iv}}~$\\lambda 1548,1550$ doublet as well as %Ly$\\alpha$ and one shows the {\\ion{Mg}{ii}}~$\\lambda\\lambda 2796,2803$ %doublet and the {\\ion{Mg}{i}}~$\\lambda 2852$ line. } % conclusions heading (optional), leave it empty if necessary {} %The redshift of GRB\\,090926A is high enough to allow us the detection %  of Ly$\\alpha$ and Ly$\\beta$ in absorption for the host galaxy, and %  to compute %  The Hydrogen column density associated to GRB\\,090926A is $\\log %  N_{\\rm H}/{\\rm cm}^{-2} = 21.60 \\pm 0.07$, and the metallicity of the %  host galaxy is in the range [X/H] $= %  -2.5$ to $-1.9$ with respect to the solar values,   %  i.e., among the lowest values ever observed for a GRB host galaxy. A %  comparison with galactic chemical evolution models has suggested %  that the host of GRB\\,090926A is likely to be a dwarf irregular %  galaxy. No emission lines were detected, but a %  H$\\alpha$ flux in emission of $9\\times10^{-18}$ %  erg~s$^{-1}$~cm$^{-2}$ (i.e., a star formation rate of %  $2~M_\\odot$yr$^{-1}$), which is typical of many GRB hosts, would %  have been detected in our spectra, and thus emission lines are well %  within the reach of X-shooter. %%The non-detection of molecular %%  absorption, and in particular of the H$_2$ molecule, allows us to %  We put an upper limit to the \\ion{H} molecular fraction of the host %  galaxy ISM, which is $f < 7\\times10^{-7}$.  The continuum has been %  fitted assuming a power-law spectrum, with a spectral index of %  $\\beta = 0.89^{+0.02}_{-0.02}$. The best fit does not essentially %  require any intrinsic extinction since $E_{B-V} < 0.01$ mag %  adopting a SMC extinction curve.  %    We derive information on the distance between the host absorbing %    gas and the site of the GRB explosion. %%We do so by assuming that %%    the excited absorption lines are produced by indirect UV pumping, %%    and compare the data with a time dependent photo-excitation code. %    The distance of component I is found to be $2.40 \\pm 0.15$ kpc, %    while component II is located far away from the GRB, possibly at %    $\\sim 5$ kpc. These values are compatible with that found for %    other GRBs.  %    One intervening system only features {\\ion{Mg}{ii}} %    absorption, and the rest frame equivalent width of the %    {\\ion{Mg}{ii}}~$\\lambda 2796$ is smaller than 1~\\AA. The %    GRB\\,090926A line of sight is thus much cleaner than the line %    of sights to other GRBs, where one strong intervening system is %    typically found per afterglow, as has been noted in studies %    showing excess of strong {\\ion{Mg}{ii}} absorbers along GRB %    sightlines compared to quasars. % %} ",
        "introduction": "For a few hours after their onset, long gamma-ray burst (GRB) afterglows (see, e.g., M\\'esz\\'aros 2006 for a review) are the brightest sources in the Universe. Their intense and featureless continua in the X-ray to the radio band offer a superb opportunity to study the gas and dust content of high redshift galaxies (see e.g., Vreeswijk et al. 2001; Vreeswijk et al. 2004; Berger et al. 2006; Fynbo et al. 2006; Prochaska et al. 2007; Sparre et al. 2013). Therefore, GRBs provide a complementary diagnostic tool to study high-redshift environments, with respect to Lyman-break galaxies (LBGs, see e.g. Steidel et al.  1999), $K$-band selected galaxies (Savaglio et al. 2004; Martin et al. 2012; Kornei et al. 2013), and galaxies that happen to be along the lines of sight to bright background QSOs. In fact, LBG studies are biased towards high UV-luminosity galaxies, and their inferred chemical abundances may not be related to that of typical high-redshift galaxies.  Ultra-deep spectroscopic observations may allow us to study normal galaxies, but only up to $z \\sim 2$ (using the present generation of 8-m class telescopes). For example, Savaglio et al. (2004, 2005) studied the inter-stella medium (ISM) of a sample of faint $K$-band selected galaxies, finding abundances much higher than in QSOs systems. This is because the line of sight to QSO probes mainly galaxy halos, rather than their bulges or discs (Fynbo et al. 2008), returning lower abundance values. GRBs represent a way to study the central regions of high-redshift galaxies, without luminosity biases. Abundances measured in GRB host galaxies vary from less than $10^{-2}$ to solar values and are on average higher than those found along QSO sightlines (see e.g., Fynbo et al. 2006; Prochaska et al. 2007; Savaglio et al. 2012; Sparre et al. 2013). GRBs are expected to originate in molecular clouds, where star formation occurs. Nevertheless, absorption from ground-state and vibrationally excited levels of the H$_2$ and other molecules is hardly observed in GRB afterglow spectra (Vreeswijk et al. 2004, Tumlinson et al. 2007), with the exception of GRB\\,080607 (Prochaska et al. 2009, Sheffer et al. 2009), GRB\\,120815A (Kr\\\"uhler et al. 2013), and possibly GRB\\,060206 (Fynbo et al. 2006). The non-detection of molecules is possibly a consequence of photo-dissociation by the intense UV flux from the GRB afterglow or by the young, hot stellar populations harboring the GRB host.  Furthermore, GRBs are the most variable and violent phenomena in the Universe, while QSOs are persistent, albeit variable, sources. Thus, the physical, dynamical and chemical status of the circumburst medium in the star-forming region hosting GRB progenitors is altered in real time by the explosive event, through shock waves and ionizing photons. The transient nature of GRBs is manifested through variability in the populations of fine structure and other excited levels of the {\\ion{O}{i}} atom and of ions such as {\\ion{Fe}{ii}}, {\\ion{Ni}{ii}}, {\\ion{Si}{ii}} and {\\ion{C}{ii}}. The population of these excited levels, and therefore the observed column densities of the corresponding absorption lines, change on timescales of minutes to days (Vreeswijk et al. 2007, D'Elia et al. 2009a, De Cia et al. 2012, Hartoog et al. 2013). This variation is not consistent with a pure infrared excitation or collisional processes, but is due to UV-pumping (Prochaska, Chen \\& Bloom 2006, Vreeswijk et al. 2007). Monitoring these variations and comparing them with predictions from photo-excitation codes enables one to derive information on the distance between the GRB and the absorbing material. This distance can be estimated even when multi-epoch spectroscopy is not available, under the assumption that UV pumping is the responsible mechanism for the production of the excited lines (D'Elia et al. 2009b; 2010; 2011; Ledoux et al. 2009).  %GRB spectroscopy is also suitable for studying systems lying along the %line of sight to GRBs.  Surprisingly, the number density of strong %{\\ion{Mg}{ii}} intervening absorbers is more than twice higher than %along QSO sightlines, considering high resolution spectra only %(Prochter et al. 2006, Vergani et al. 2009), while {\\ion{C}{iv}} %absorbers do not show any statistical difference (Sudilovsky et %al. 2007, Tejos et al. 2007). The reason for this {\\ion{Mg}{ii}} %excess in GRB high resolution spectra is still unclear (Porciani et %al. 2007, Cucchiara et al. 2009), and puzzling, because this %discrepancy seems to disappear when considering also low resolution %spectra (Cucchiara et al. 2013).  %All these issues can now be systematically addressed using the %X-shooter spectrograph. This is the first second-generation instrument %at the ESO's Very Large Telescope (VLT) at Paranal Observatory %(Chile). It is a single target-spectrograph capable of obtaining a %medium-resolution spectrum ($R=\\lambda/\\Delta\\lambda=4000 - 14,000$) %covering the spectral range 3000 - 24,800~\\AA{} in a single exposure %thanks to splitting the light into three arms: ultraviolet/blue (UVB), %visual (VIS), and near-infrared (NIR). We refer the reader to %D'Odorico et al. (2006) for a complete description of the X-shooter %specifications and capabilities, and to de Ugarte Postigo et %al. (2010) for the first results on the GRB field.  The GRB environments are in general very complex. In fact, they are constituted by several layers of gas which may be close to or far from one another, have different relative velocities and physical properties. Optical spectroscopy reflects this complexity, since the absorption lines in several cases cannot be fitted with a single line profile. The ISM complexity is particularly evident for GRB afterglows observed at high resolution (Fiore et al. 2005; D'Elia et al. 2007; Th\\\"one et al. 2008; Piranomonte et al. 2008).  Here we present the case of GRB\\,120327A, observed by X-shooter on 2012 March 27 and 28. Our dataset is complemented by a GTC spectrum and photometric data from several facilities. We derive the abundances by comparing the column densities of hydrogen and metals, and evaluate the GRB-absorber distance using the variation of the excited levels and their ratios to the ground state ones, and the predictions of photo-excitation codes. We also characterize the host galaxy using the dust depletion pattern technique and chemical evolution models, derive the extinction curve shape and search for H$_2$ molecules in the GRB environment.  The paper is organized as follows. Section 2 summarizes the properties of GRB\\,120327A and results of observations previously reported in the literature; Sect. 3 presents our observations and data reduction; Sect. 4 is devoted to the analysis of the absorption features from the host galaxy, using X-shooter data; Sect. 5 illustrates our main results, which are discussed in Sect. 6. Finally, in Sect. 7 the conclusions are drawn. We assume a cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, $\\Omega_\\Lambda = 0.7$. Hereafter, the notation [X/H] indicates the logarithmic abundance of element X relative to solar values. ",
        "conclusions": "In this paper we present multi-epoch, intermediate-resolution ($R=7000 - 12000$) spectroscopy of the optical afterglow of GRB\\,120327A, observed using the X-shooter spectrograph at the VLT, $\\sim 2.5$ and $27.7$ hr after the trigger.  Our main findings are summarized below.  (i) From the detection of hydrogen and metal absorption features, we find that the redshift of the host galaxy is $z=2.8145\\pm0.0001$. The ISM of the GRB host galaxy has at least three components contributing to the absorption. The line center of component I (III) is $41$ ($35$) km s$^{-1}$ bluewards (redwards) of component II, which sets the redshift reference point.  (ii) Component III presents less absorption features and no excited lines, and is probably the farthest from the GRB explosion site.  (iii) Components I and II have similar absorption profiles, but also some differences, with the former showing more absorption from neutral elements and the latter from highly ionized species. Fine structure and metastable levels are present in both components, but with higher column densities in component II. This points towards component II being closer to the GRB than component I.  (iv) High ionization lines do not present blue wings or higher velocity components with respect to low ionization ones, but show a similar velocity range, contrary to what is observed in previous GRBs.  (v) The distance of these components from the GRB is computed using a time dependent photo-excitation code and comparing the results with the excited line ratios for the \\ion{Fe}{ii} ion. Component I is $200^{+100}_{-60}$ pc away from the GRB explosion site, while component II is at $100^{+40}_{-30}$ pc, confirming result (iii). The variability of the \\ion{Fe}{ii}~$\\lambda$2396 fine structure line between the two X-shooter spectra proves that the excited levels are produced by the GRB UV flux.  (vi) The column density of the \\ion{Fe}{ii} ground state is in excess from what is expected using the photo-excitation code, both in component I and II. This possibly indicates that most of \\ion{Fe}{ii} lies far away from the GRB, and that the geometry of the GRB\\,120327A host galaxy absorber is more complex with respect to the resolving capacity of X-shooter.  (vii) The damped Lyman-$\\alpha$ system associated with the GRB host galaxy is prominent, with a column density of $\\log (N_{\\rm H}/{\\rm cm}^{-2}) = 22.01 \\pm 0.09$. The metallicity computed using the less depleted elements is between $\\mbox{[X/H]} = -1.3$ and $-1.1$, i.e., a value in the middle of the metallicity distribution for the GRB host galaxies at high redshift.  (viii) The dust depletion pattern indicates an $\\alpha$ element enhancement with respect to the elements of the Iron group, typical of massive galaxies. The same conclusion is supported by modeling this overabundance using detailed chemical enrichment codes.  (ix) A joint analysis of host galaxy photometry and chemical enrichment codes leads us to suggest that the host is a high mass ($>10^{10}$ $M_{\\odot}$), high luminosity one.  (x) Concerning the dust content in the host, we estimate $A_V = 0.22$ mag along the line of sight of GRB 120327A, using the dust depletion pattern, considerably in excess of what we find studying the extinction curve shape by modeling the flux-calibrated spectra ($A_V < 0.03$ mag). The ratio $\\mbox{[Zn/Fe]} = 0.56\\pm0.14$ depicts an environment with a certain dust depletion, but not so extreme as in other dusty GRBs such as GRB\\,120815A or GRB\\,080607.  (xi) The H$_2$ molecule is firmly detected in the GRB\\,120327A host galaxy. This is the third secure detection of molecules in a GRB environment after GRB\\,080607 and GRB\\,120815A. The molecular fraction we determined is at least two orders of magnitude below that of the other GRBs, possibly because of the lower dust content, since it helps the molecule formation."
    },
    "1402/1402.0526.txt": {
        "abstract": "I give a general review of inflationary cosmology and of its present status, in view of the 2013 data release by the Planck satellite. A specific emphasis is given to the new broad class of theories, the cosmological attractors, which have nearly model-independent predictions converging at the sweet spot of the Planck data in the  $(n_{s},r)$ plane. I also discuss the problem of initial conditions for the theories favored by the Planck data. ",
        "introduction": "We were waiting for the recent Planck data release for many years. There were persistent rumors that this data are going to confirm the existence of relatively large non-Gaussian perturbations, which would rule out the simplest single-field inflationary models, but would open the door towards complicated multi-field models of inflation, which would provide lots of work for many generations of cosmologists. For several years, many scientists were preparing themselves to this grave eventuality. But instead of that, Planck 2013 found that the non-Gaussian perturbations of the most dangerous type are practically absent \\cite{Ade:2013uln}, and therefore one could return to investigation of the simplest and nicest inflationary models. Some of us, me included, were ecstatic to hear the news, but for many others it was a disappointment. It could seem that now we have nothing new to say about inflation since everything related to its simplest versions was already known, classified, forgotten, and rediscovered over and over again during the last 30 years. However, as I will argue in these lectures, the new results obtained by Planck pushed us in an unexpected direction, which stimulated development of new concepts and new broad classes of inflationary models with most interesting properties. In addition to that, the split between the inflationary cosmology and its alternatives became even more apparent after the Planck 2013. Whereas the Planck team emphasized that their results confirm basic principles of inflation and rule out many alternative theories \\cite{Ade:2013uln}, some of the proponents of these alternatives arrived at an absolutely opposite conclusion. Thus life after Planck 2013 became quite exciting. Hopefully these lectures may help to understand the essence of the recent developments and debates. But in order to do it, we must start from the very beginning. For the first chapters, I will borrow some materials from my lectures written several years ago \\cite{Linde:2007fr}, and then we will fast forward to the very recent developments. Finally, I should note that some issues closely related to the subjects covered in my lectures have been discussed by other speakers at the Les Houches school. In particular, the approach to cosmology based on supergravity and superconformal theory was described by Renata Kallosh \\cite{RenataLecture}, and string theory approach to  inflation was presented by Eva Silverstein \\cite{Silverstein:2013wua}.   \\section {Brief history of inflation}  Several ingredients of inflationary cosmology were discovered in the beginning of the 70's. The first realization was that the energy density of a scalar field plays the role of the vacuum energy/cosmological constant \\cite{Linde:1974at}, which was changing during the cosmological phase transitions \\cite{Kirzhnits:1972ut}. In certain cases these changes  occur discontinuously, due to first order phase transitions from a supercooled vacuum state (false vacuum) \\cite{Kirzhnits:1976ts}.  In 1978, Gennady Chibisov and I tried to use these facts to construct a cosmological model involving exponential expansion of the universe in a supercooled vacuum as a source of the entropy of the universe, but we immediately realized that the universe becomes very inhomogeneous after the bubble wall collisions. I mentioned our work in my review article \\cite{Linde:1978px}, but did not pursue this idea any further.   The first semi-realistic model of inflationary type was proposed by Alexei Starobinsky in 1980 \\cite{Starobinsky:1980te}. It was based on investigation of conformal anomalies in quantum gravity.  His model was rather complicated, and its goals were in a certain sense opposite to the goals of inflationary cosmology. Instead of attempting to solve the homogeneity and isotropy problems, Starobinsky considered the model of  the universe which was homogeneous and isotropic from the very beginning, and emphasized that his scenario was ``the extreme opposite of Misner's initial chaos''   \\cite{Starobinsky:1980te}.  On the other hand, Starobinsky's model did not suffer from the graceful exit problem, and it was the first model predicting  gravitational waves with a flat spectrum \\cite{Starobinsky:1980te}. The first mechanism of production of adiabatic perturbations of the metric with a nearly flat spectrum, which are responsible for galaxy production, and which were found by the observations of the CMB anisotropy,  was proposed by  Mukhanov  and Chibisov \\cite{Mukh} in the context of this model. They also pointed out that the spectrum should not be exactly flat, the fact that was confirmed by Planck 2013 at the 5$\\sigma$ level.   A much simpler inflationary model with a very clear physical motivation was proposed by Alan Guth in 1981 \\cite{Guth}.  His model, which is now called ``old inflation,'' was based on the theory of supercooling during the cosmological phase transitions \\cite{Kirzhnits:1976ts}. Even though this scenario did not work in its original form,  it played a profound role in the development of inflationary cosmology since it contained a very clear explanation how inflation may solve the major cosmological problems.  According to this scenario,  inflation is as   exponential expansion of the universe in a supercooled false vacuum state. False vacuum is a metastable state without any fields or particles but with large energy density. Imagine a universe filled with such ``heavy nothing.'' When the universe expands, empty space remains empty, so its energy density does not change. The universe with a constant energy density expands exponentially, thus we have inflation in the false vacuum. This expansion makes the universe very big and very flat. Then the false vacuum decays, the bubbles of the new phase collide, and our universe becomes hot.  Unfortunately, this simple and intuitive picture of inflation in the false vacuum state is somewhat misleading. If the probability of the bubble formation is large, bubbles of the new phase are formed near each other, inflation is too short to solve any problems, and the bubble wall collisions make the universe extremely inhomogeneous.  If they are formed far away from each other, which is the case if the probability of their formation is small and inflation is long, each of these bubbles represents a separate open universe with a vanishingly small $\\Omega$. Both options are unacceptable, which has lead to the conclusion that this scenario does not work and cannot be improved (graceful exit problem) \\cite{Guth,Hawking:1982ga,Guth:1982pn}.   A solution of this problem was found in 1981 with the invention of the new inflationary theory \\cite{New}. (Three months later, the same solution was also proposed in \\cite{New2}, with a reference to my earlier work \\cite{New}.) In this theory, inflation may begin either in the false vacuum,  or in an unstable state at the top of the effective potential. Then the inflaton field $\\phi$ slowly rolls down to the minimum of its effective potential.  The motion of the field away from the false vacuum is of crucial importance: density perturbations produced during the slow-roll inflation are inversely proportional to $\\dot \\phi$ \\cite{Mukh,Hawk,Mukh2}. Thus the key difference between the new inflationary scenario and the old one is that the useful part of inflation in the new scenario, which is responsible for the homogeneity of our universe, does {\\it not} occur in the false vacuum state, where $\\dot\\phi =0$.   Soon after the invention of the new inflationary scenario it became very popular. Unfortunately, this scenario was plagued by its own problems. In most versions of this scenario the inflaton field must have an extremely small coupling constant, so it could not be in thermal equilibrium with other matter fields, as required in  \\cite{Guth,New}. The theory of cosmological phase transitions, which was the basis for old and new inflation, did not work in this case. Moreover, thermal equilibrium requires many particles interacting with each other. This means that new inflation could explain why our universe was so large only if it was very large and contained many particles from the very beginning \\cite{Linde:2005ht}.  Old and new inflation represented a substantial but incomplete modification of the big bang theory. It was still assumed that the universe was in a state of thermal equilibrium from the very beginning, that it was relatively homogeneous and large enough to survive until the beginning of inflation, and that the stage of inflation was just an intermediate stage of the evolution of the universe. In the beginning of the 80's these assumptions seemed most natural and practically unavoidable. On the basis of all available observations (CMB, abundance of light elements) everybody believed that the universe was created in a hot big bang. That is why it was so difficult to overcome a certain psychological barrier and abandon all of these assumptions. This was done in 1983 with the invention of the chaotic inflation scenario \\cite{Linde:1983gd}. This scenario resolved all problems of old and new inflation. According to this scenario, inflation may begin even if there was no thermal equilibrium in the early universe, and  it may occur even in the theories with simplest potentials such as $V(\\phi) \\sim \\phi^2$. But it is not limited to the theories with polynomial potentials: chaotic inflation occurs in {\\it any} theory where the potential has a sufficiently flat region, which allows the existence of the slow-roll regime \\cite{Linde:1983gd}.  This change of perspective happened 30 years ago, and since that time cosmologists almost never returned to the original idea of inflation due to high temperature phase transitions in grand unified theories. This idea was abandoned by all experts actively working in this area of physics, but unfortunately it is still present in virtually all textbooks on astrophysics, which incorrectly describe inflation as an exponential expansion in a supercooled false vacuum state during the cosmological phase transitions in grand unified theories. Apparently the front lines of the development of this branch of science moved forward too fast, so let us remember what exactly happened with the development of chaotic inflation. ",
        "conclusions": "Three decades ago, inflationary theory looked like an exotic product of vivid scientific imagination. Some of us believed that it possesses such a great explanatory potential that it must be correct; some others thought that it is too good to be true. Not many expected that it is possible to verify any of its predictions in our lifetime.  Thanks to the enthusiastic work of many scientists, inflationary theory is gradually becoming a broadly accepted cosmological paradigm, with many of its predictions being confirmed by observational data.  The new data release by Planck 2013 stimulating the development of new cosmological theories, by changing the goal from finding various complicated models capable of describing large local non-Gaussiantiy to the development of new elegant models of inflation capable of explaining increasingly precise data in the $(n_{s},r)$ plane. The simplicity of the chaotic inflation scenario with potentials $\\phi^{n}$ is unmatched, but what if the observational data force us to modify these models? Just few months ago, we did not know many good inflationary models which would naturally predict the data favored by Planck 2013. Now the situation has changed. The existence of the universal attractor regime for a large set of different inflationary models does not guarantee that we are on the right track, but it is hard to ignore that all of these cosmological attractors point in the same direction, and their predictions converge at the ``sweet spot'' in the $(n_{s},r)$ plane preferred by WMAP9 and Planck 2013. It may be equally important that if one follows the attractor trajectories in the opposite direction (i.e. takes the limit $\\xi \\to 0$ or $\\alpha \\to \\infty$ in this class of models), one continuously interpolates back to the predictions of the simplest models of chaotic inflation.   %\\subsubsection*"
    },
    "1402/1402.1214_arXiv.txt": {
        "abstract": "We present a complete catalog of \\HI emission and absorption spectrum pairs, toward \\HII regions, detectable within the boundaries of the Southern Galactic Plane Survey (SGPS I), a total of 252 regions.  The catalog is presented in graphical, numerical and summary formats.  We demonstrate an application of this new dataset through an investigation of the locus of the Near 3kpc Arm. ",
        "introduction": "Combining data from the Australia Telescope Compact Array (ATCA) and the Parkes single dish telescopes, the SGPS  \\citep[SGPS,][]{SGPS} provides \\HI line and 1.4 GHz radio continuum data for the fourth Galactic quadrant, with the best combination of resolution and sensitivity currently available for that line in this area.  Using these data, we have measured 21-cm absorption spectra toward a sample of 252 \\HII regions.  This study produces two distinct data products: a consolidated census of \\HII regions with known radio recombination line (RRL) velocities, taken from the literature; as well as the \\HI emission and absorption spectrum pairs towards them - within the bounds of the SGPS (255\\arcdeg$<l<$353\\arcdeg).  These two resources will enable further study into the structure and dynamics of the neutral interstellar medium (ISM) in the fourth quadrant (on which massive stars and their formation have a significant impact).  The catalogs will serve as a data set for numerous studies; including investigations of the spiral structure of the Galaxy \\citep[eg. ][]{Strasser07} and kinematic distance works, both for target regions \\citep[eg.][]{Urquhart12} as well as for the intervening \\HI clouds \\citep[eg.][]{RomanDuval09}.  The sample selection from existing \\HII region catalogs is described in Section \\ref{Data}, while the method of spectrum extraction appears in Section \\ref{extract}.  The catalog itself appears in Section \\ref{Catalog}, before a disucssion of the global properties of \\HI absorption in Section \\ref{Discussion} and an illustration of its use: an examination of the distribution of \\HI absorption in longitude-velocity ($lv$) space, including the locus of the Near 3kpc Arm. ",
        "conclusions": "\\label{Discussion} Here we briefly discuss the global properties of the \\HI absorption catalog (\\S \\ref{disscusslv}), as well as an example of its use - investigating the $lv$ locus of the Near 3kpc Arm (see \\S \\ref{N3kpc}).  \\subsection{Longitude-Velocity Distribution}\\label{disscusslv} One simple application of this catalog is an investigation of the distribution of \\HI absorption in longitude-velocity ($lv$) space.  Firstly, we construct an `$lv$ crayon diagram\\rq{} (Figure \\ref{Fig5}) of known Galactic structures.  The Near and Far 3kpc Arm fits are provided by \\citet{DameThaddeus08} and the fourth quadrant spiral arms are taken from \\citet{Vallee08} (who represents the Distant Arm of \\citet{McCGDistantArm} as the Cygnus Arm beyond the Solar Circle).  The velocity width of each crayon feature is set to 20\\kms.  For each source in the catalog (of quality factor C or better), we plot both the systemic velocity of the \\HII region and any associated \\HI absorption (Figure \\ref{Fig5}).  It is clear that both the \\HII regions and their associated \\HI absorption trace the spiral arm structures, especially in the inner Galaxy.  This is in keeping with \\citet{Jones13} who find that $>90\\%$ of their \\HII region sample is associated with known spiral arm structures.  \\HII regions are considered to be the archetypical tracers of Galactic spiral structure \\citep{GBTHRDS}; but cold, dense gas, traced by \\HI absorption is also more likely to be located within the spiral arms than in the inter-arm region.  Each black marker on the Figure corresponds to five consecutive velocity channels which display significant absorption.  It is clear that the \\HI absorption is also associated with the spiral arms.  This is to be expected, as dense, cold gas which is traced by \\HI absorption is more likely to be located within the spiral arms than in the inter-arm regions.  We find \\HI absorption associated with all inner Galaxy spiral arms.  Figure \\ref{Fig5} not only demonstrates the suitability of both data-sets (\\HII region velocities and \\HI absorption) as spiral arm tracers, but also their inherent velocity uncertainty distributions.  Galactic streaming motions are often estimated to be on the order of $\\lesssim15$\\kms \\citep{BaniaLockman84, Kolpak03}, i.e. within the velocity width of the crayon diagram features.  However, \\HI absorption associated with \\HII regions can extend $\\sim20$\\kms beyond the systemic velocity of the region \\citep{Dickey03, Jones13}.  The association of \\HI absorption with a Galactic structure feature can be used to infer lower distance limits, or the location of the \\HII region in our Galaxy \\citep{JonesDickey12, Jones13}.    \\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{Figure5} \\caption{\\HII region RRL velocities (red crosses) and \\HI absorption (black dots) overlaid on a longitude-velocity `crayon diagram'.  Each dot denotes the central velocity of five adjacent velocity channels ($\\sim5$\\kms), with each channel demonstrating a statistically significant amount of \\HI absorption.  Only \\HII regions with \\HI absorption spectra of Quality 'C' or better are included for use in this figure. (A color version of this figure is available in the online journal.)}\\label{Fig5} \\end{figure*}  \\subsubsection{\\HI Absorption in the Outer Galaxy} From Figure \\ref{Fig5} it is evident that this catalog is primarily limited to the inner Galaxy (within the Solar circle); unlike, for example, the work of \\citet{Strasser07}.  \\citet{Strasser07} used the SGPS and the E/A spectrum extraction method to study \\HI absorption towards 111 extragalactic continuum sources - in order to investigate absorption from the outermost arms of the Milky Way.  Of these 111 sources, only 17 demonstrated any signature of \\HI absorption at positive velocities; corresponding to a location in the outer Galaxy.  Because the \\citet{Strasser07} source sample included strictly extra-Galactic sources, the \\HI absorption spectra reflect lines of sight through the entire Galactic plane.  Whereas spectra from this study reflect lines of sight towards the \\HII regions themselves, located within the Galaxy.  For this reason, this work is far less likely to identify \\HI absorption in the outermost spiral arms in the fourth quadrant - as the number of \\HII regions known (with RRL velocities) is extremely limited in the most distant features.  Nevertheless, for longitudes $l<300\\arcdeg$ the line of sight distance to the Solar circle is relatively small and investigations of \\HI absorption in the outer Galaxy become possible; although at these longitudes, there are very few known \\HII regions (see Figure \\ref{Fig5}).  Regions with velocities $\\lesssim15$\\kms correspond to very short line of sight distances.  At $l\\approx 250\\arcdeg$ the near-continuous absorption profile is associated with lines of sight through the \\objectname[LOCAL ARM]{Local Arm}.  \\subsection{Locus of the Near 3kpc Arm \\label{N3kpc}}  \\begin{figure}[h!] \\centering \\includegraphics[width=\\columnwidth]{Figure6} \\caption{\\HI absorption associated with the Near 3kpc Arm.  \\textit{Left Panel}: $lv$ diagram of \\HI absorption corresponding to the Near 3kpc Arm.  Light grey dots are data from \\citet{Jones13}, dark grey dots are from this work (but were not included in the linear fit analysis), black dots are also from this work (and were used in the analysis of the locus of the Arm). \\textit{Right Panel:} fit of data to model (Equation \\ref{Eqboth}).  The histogram displays the difference in velocity between Equation \\ref{Eqboth} and each channel that displays significant absorption (see Figure \\ref{Fig5}).  Note only regions with $l>340\\arcdeg$ are considered here.\\label{Fig6}} \\end{figure} The investigation of the $lv$ distribution of \\HI absorption (see Figure \\ref{Fig5}) can be further specialized to individual features.  In \\citet{Jones13} the locus of the Near and Far 3kpc Arms (in $lv$ space) was investigated using \\HI absorption.  This work presents an opportunity to extend this analysis, by including \\HI absorption associated with the Near 3kpc Arm at longitudes $l<350\\arcdeg$.  This is a complementary investigation to that first performed by \\citet{DameThaddeus08} using CO 1-0 emission.  Using \\HI absorption channels from the Near 3kpc Arm analysis of \\citet{Jones13} in addition to those from this work---such that the entire longitude extent 10\\arcdeg$>l>340$\\arcdeg of the Near 3kpc Arm is included---a linear fit to the $lv$ locus of the Near 3kpc Arm was performed: \\begin{equation} V=-53.3+4.1l\\qquad\\Delta V=25.9\\textrm{\\kms} \\label{Eqboth} \\end{equation} Here $\\Delta V$ refers to the FWHM of a Gaussian fit to the velocity profile of the Arm model (see right hand panel of Figure \\ref{Fig6}).  With the inclusion of \\HI absorption from the longitude range $l<340\\arcdeg$, the model for the locus of the Near 3kpc Arm (Equation \\ref{Eqboth}) is extremely consistent with the fit provided by \\citet{DameThaddeus08} from CO observations; $R^2>0.97$ and the standard error of the mean difference is 5.4\\kms, within the FWHM of either fit.  Figure \\ref{Fig6} demonstrates the location of \\HI absorption associated with the Arm in $lv$ space, as well as an evaluation of the fit.  The left hand panel,  supplementary to Figure \\ref{Fig5}, is an $lv$ diagram extended to $l=10\\arcdeg$ over the velocity range of the Near 3kpc Arm.  Channels of significant absorption are shown,  in addition to those from \\citet{Jones13}.  The $lv$ linear fits to the locus of the Near 3kpc Arm from \\citet[][$V=-53.1+4.6, \\Delta V=19.7$ \\kms]{DameThaddeus08}, \\citet[][$V=-59.2+4.12, \\Delta V=17.3$ \\kms]{Jones13} and this work (see Equation \\ref{Eqboth}) are also shown.  In the right hand panel an evaluation of the fit is made.  All velocity channels shown in the left-hand panel are compared to the linear Near 3kpc Arm fit  (similar to the CO 1-0, Figure 2 of \\citet{DameThaddeus08}).  The Near 3kpc Arm is identified as a well defined peak in \\HI absorption, centered around the fit to the linear model (however, the scale height of the component is exaggerated by the removal of non-Near 3kpc Arm absorption for $360\\arcdeg>l>353\\arcdeg$).  The hump at $\\Delta V=30$ (i.e. velocities consistent with a line parallel to Equation \\ref{Eqboth} $+30$\\kms) may be the signature of the Norma Arm, which is approximately parallel to the Near 3kpc Arm in $lv$ space for longitudes $l>340\\arcdeg$ (see Figure \\ref{Fig5}).  The broad central peak is the result of the remaining disk rotation, both foreground and background to the Near 3kpc Arm.  Note that longitudes $l<350\\arcdeg$ the Far 3kpc Arm is confused in velocity with the other inner Galaxy spiral arms - therefore investigation of the locus of the Far 3kpc Arm, extending the work of \\citet{Jones13}, is not possible here."
    },
    "1402/1402.6328_arXiv.txt": {
        "abstract": "We present $\\mathrm{\\alpha}$-element abundances of Mg, Si, and Ti for a large sample of field stars in two outer fields of the Fornax dwarf spheroidal galaxy (dSph), obtained with VLT/GIRAFFE ($R\\sim16,000$). Due to the large fraction of metal-poor stars in our sample, we are able to follow the $\\mathrm{\\alpha}$-element evolution from $\\mathrm{[Fe/H]}\\approx-2.5$ continuously to $\\mathrm{[Fe/H]}\\approx-0.7$. For the first time we are able to resolve the turnover from the Type II supernovae (SNe) dominated, $\\mathrm{\\alpha}$-enhanced plateau down to subsolar $\\mathrm{[\\alpha/Fe]}$ values due to the onset of SNe Ia, and thus to trace the chemical enrichment efficiency of the galaxy. Our data support the general concept of an $\\mathrm{\\alpha}$-enhanced plateau at early epochs, followed by a well-defined ``knee'', caused by the onset of SNe Ia, and finally a second plateau with sub-solar $\\mathrm{[\\alpha/Fe]}$ values. We find the position of this knee to be at $\\mathrm{[Fe/H]}\\approx-1.9$ and therefore significantly more metal-poor than expected from comparison with other dSphs and standard evolutionary models. Surprisingly, this value is rather comparable to the knee in Sculptor, a dSph $\\sim10$ times less luminous than Fornax. Using chemical evolution models, we find that both the position of the knee as well as the subsequent plateau at sub-solar level can hardly be explained unless the galaxy experienced several discrete star formation events with a drastic variation in star formation efficiency, while a uniform star formation can be ruled out. One possible evolutionary scenario is that Fornax experienced one or several major accretion events from gas-rich systems in the past, so that its current stellar mass is not indicative of the chemical evolution environment at ancient times. If Fornax is the product of several smaller buildings blocks, this may also have implications of the understanding on the formation process of dSphs in general. ",
        "introduction": "\\label{chap_01}  Dwarf spheroidal (dSph) galaxies in the Local Group are the smallest, closest and most abundant galaxies in the Universe and are therefore excellent laboratories to study galactic star formation history (SFH) and chemical evolution on the smallest scale. Among these, Fornax is one of the most massive (stellar mass $\\sim10^7 M_{\\odot}$, \\citealt{de_Boer_12}) and most metal-rich ($\\langle\\mathrm{[Fe/H]}\\rangle \\approx-1.0 $) dSphs at a distance of $147\\,\\mathrm{kpc}$ (\\citealt{McConnachie_12}), and one of only two systems known to host a globular cluster (GC) population (the other one is Sagittarius). Fornax displays complex chemical and dynamical properties: the metallicity-distribution-function (MDF) from several spectroscopic studies (\\citealt{Battaglia_06}, Hendricks et al. in prep) shows distinct peaks in [Fe/H] around $-0.5$, $-1.0$, $-1.4$ and $-2.0$\\,dex, which can be interpreted as an inhomogeneous, bursty SFH. In contrast, photometric SFH studies find continuous star formation (SF) from a population older than 10--12\\,Gyr continuing to a young population not older than several hundred Myr with a significant SF peak around 3--4\\,Gyr ago (\\citealt{de_Boer_12}, \\citealt{Coleman_08}). The age resolution of photometric studies, however, is limited due to the low color-sensitivity of old RGB stars and additionally biased by systematic uncertainties in the Calcium Triplet (CaT) metallicites used to break the age-metallicity degeneracy (\\citealt{Battaglia_08}).  Possible proposed origins for the complex chemical evolution of Fornax are periodic interactions of the galaxy with the MW's tidal field (e.g., \\citealt{Nichols_12}, see also \\citealt{Pasetto_11} for Carina), or merger scenarios with other dwarf galaxies (\\citealt{Coleman_04}, \\citealt{Yozin_12}). On the other hand, potential SF bursts may also be a consequence of the reaccretion of gas previously blown out of the galaxy by stellar winds with a time delay of several $10^9$\\,years (e.g., \\citealt{Revaz_09}, \\citealt{Ruiz_13}). Interestingly, two of the observed peaks in the field star MDF roughly coincide with the metallicity of Fornax' GCs and it has been proposed that these peaks are not caused by intrinsic star formation variations, but rather by dissolved GC stars (\\citealt{Larsen_12}). These arguments show that the SFH of Fornax is not well understood yet, and an important question remaining is whether the system evolved in unperturbed isolation, or whether its field population is the result of merging events and/or tidal interactions with the MW.  The $\\mathrm{\\alpha}$-elements (e.g., Mg, Si, Ca, Ti) provide a crucial tool for understanding a galaxy's chemical enrichment history. They are synthesized mainly in massive stars ($M_{\\star} \\geq 8M_{\\sun}$) with a lifetime of not more than 30\\,Myr, and distributed throughout a galaxy during the subsequent core-collapse SN\\,II explosion. In contrast, SNe Ia are thought to produce primarily Fe--peak elements and may not contribute to a system's interstellar medium (ISM) for $\\sim10^{9}$ years. Therefore, the $\\mathrm{[\\alpha/Fe]}$-ratio is a tracer for the relative contribution of SNe\\,II and SNe\\,Ia to the ISM. At early times, only SN\\,II enrich the ISM which is observed as an $\\mathrm{\\alpha}$-rich plateau at very low values of [Fe/H]. The sudden drop in the $\\mathrm{[\\alpha/Fe]}$-ratio caused by the onset of SN\\,Ia is generally called the ``knee\" in the evolution of $\\mathrm{\\alpha}$-elements. Consequently, the knee in a plot of $\\mathrm{[\\alpha/Fe]}$ versus [Fe/H] is an indicator for the system's efficiency to build up and retain metals within its potential well (e.g., \\citealt{Tinsley_79}, \\citealt{Matteucci_90}, \\citealt{Venn_04}, but see \\citealt{McWilliam_13} for a different view).  It has been shown for several dSphs, that the $\\mathrm{\\alpha}$-element evolution is significantly different from the MW. Generally, their knee---if detected at all---lies at lower [Fe/H] than for MW halo field stars (see, e.g., \\citealt{Tolstoy_09}), and the individual $\\mathrm{\\alpha}$-element ratios become more depleted (e.g., \\citealt{Letarte_10}, \\citealt{Koch_08}, \\citealt{Sbordone_07}). Additionally, the position of the knee varies between individual dSphs (e.g., \\citealt{Cohen_09, Cohen_10}). This apparent variation in the chemical enrichment process has been linked mainly with the total stellar mass of the individual galaxy, where faint, low-mass dwarfs show lower enrichment efficiencies. The same simplified argument also applies for the observed difference between dwarfs and the MW (\\citealt{Matteucci_90}).  Unfortunately, for most of these galaxies, the detailed evolution of the $\\mathrm{\\alpha}$-elements is not well known yet. Until today the observed sample of metal-poor (MP) stars in dSphs with high-resolution (HR) spectroscopy is small. One reason for the lack of data in this metallicity regime is the generally small fraction of MP stars ($\\mathrm{[Fe/H]}\\leq-2.0$) in dSphs (\\citealt{Helmi_06}). Second, to optimize galaxy membership, most large surveys target their central regions, which are known to be more metal rich than the outer parts (e.g., \\citealt{Battaglia_06}, \\citealt{Koch_06}, \\citealt{Kirby_11c}). Although the few published abundances suggest that, at the metal-poor end, the $\\mathrm{\\alpha}$-elements in dSphs overlap with the metal-poor MW halo (\\citealt{Shetrone_03}, \\citealt{Tafelmeyer_10}), the only galaxy for which the position of the knee is well-defined so far is Sculptor (\\citealt{Starkenburg_13}). For Fornax, \\citet{Letarte_10} provides the only existing HR study. They determined the $\\mathrm{\\alpha}$-depleted level for metal-rich stars, but only detected very few stars with $\\mathrm{[Fe/H]}\\leq-1.3$, which did not allow for the determination of the knee, nor the chemical evolution at lower metallicities.  Here, we show the results for three $\\mathrm{\\alpha}$-elements (Mg, Si, Ti) in Fornax, determined from HR spectroscopy, covering the full range in [Fe/H]. We are able to define the position of the knee and trace the evolution of the $\\mathrm{\\alpha}$-elements from the $\\mathrm{\\alpha}$-rich, metal-poor plateau, to the $\\mathrm{\\alpha}$-depleted level at the metal-rich side of the knee. In Chapter \\ref{chap_02}, we summarize our data and the reduction steps to obtain our abundances, which are subsequently presented in Chapter \\ref{chap_03}. In Chapter \\ref{chap_4}, we compare our results to chemical evolution model predictions. Finally, in Chapter \\ref{chap_05} we will discuss our findings in regard to possible formation and evolution scenarios of Fornax. ",
        "conclusions": "\\label{chap_05}  We find a clear knee in the evolution of the $\\mathrm{\\alpha}$-elements as a function of [Fe/H]. Such a feature has been observed for field stars in the MW halo, and is also predicted by chemical evolution models as the delayed onset of SNe Ia. The fact that we find Fornax' knee to be as metal-poor as the knee in the $\\sim10$ times less massive and $\\sim0.6$\\,dex more metal-poor Sculptor dSph prompts the question, if, or to what extend the total stellar mass of these galaxies determined their general efficiency to build up heavy elements over time.  By comparing our results to chemical evolution models we find that a continuous SF cannot bring the knee in the $\\mathrm{\\alpha}$-elements into a consistent evolutionary scenario with the observed MDF in Fornax. In contrast, a bursty SFH with gaps between the individual SF episodes and significant variations between the SF efficiencies can not only explain the combination of the metal-poor knee and a metal-rich MDF, but at the same time provides an explanation for the $\\mathrm{\\alpha}$-depleted plateau in the data, which we find impossible to reproduce with a continuous SFH.  Although it is surprising that a uniform SFH in Fornax does not agree with the observations, it is not striking that the match between model and data improves when the synthetic SFH is designed more flexible by allowing time-dependent key parameters. Therefore it is important to evaluate the assumptions we made for the model with respect to their physical implications for the galaxy and with regard to the the findings from previous studies.  The assumption of major SF episodes separated by several hundred\\,Myr is not in agreement with the photometric SFH reconstruction in \\citet{de_Boer_12}, who found a continuous SF present at all ages out to a radius of 0.8\\,degree. On the other hand, MDFs constructed from large samples of CaT-measurements in \\citet{Battaglia_06} and \\citet{Coleman_08} show that the distribution of stars is not homogeneous, but peak at several---in fact three or four---metallicities. The MDF from our sample shows similar signatures at the same metallicities (see Figure\\,\\ref{fig_4})\\footnote{A detailed population analysis, motivated by the findings presented in this work, will be part of a forthcoming study.}. Additionally, \\citet{Amorisco_12} find that each of the three more metal-rich populations can be distinguished by distinct kinematic properties, with a possible counter-rotation between the components, indicating a clear dynamical distinction between the populations. Having in mind that the age resolution in photometric SFH studies gradually decreases (to several Gyr) towards older ages due to a smaller color-sensitivity of old RGB stars, together with the fact that iron abundances for stars from the CaT (as used in de Boer et al.) become systematically uncertain at lower metallicities, it is possible that an interruption of SF before the ISM is enriched to $\\mathrm{[Fe/H]}\\approx-1.5$ cannot be resolved in such studies. Finally, several discrete SF episodes have been observed in the Carina dSph (e.g., \\citealt{Monelli_03} an references therein), which gives an empirical validity for this evolutionary concept in dSphs. Therefore, we consider the assumption of distinct SF bursts in Fornax to be reasonable, and the implementation of three episodes in good agreement with previous findings.  The values for the best fitting SF and wind efficiency parameters are of comparable size to the corresponding best fitting parameters in other dSphs; generally the SF efficiency is low and the galactic wind intense compared to MW field stars. While the galactic wind for our model is on the lower end of what has been found for other dwarfs (there: $6.0 \\leq \\omega_{i} \\leq 13.0$), the individual SF efficiencies for Fornax lie well in the middle of the typical parameter space in other studies, where best fitting values range between $\\nu=0.03\\,Gyr^{-1}$ for Draco and $\\nu=3.0\\,Gyr^{-1}$ for the massive Sagittarius dwarf. Sculptor, in comparison, is best fitted with $(\\nu,\\omega_{i})=(0.2,13.0)$.  However, the necessity of a strong variation of the SF efficiency (by a factor of $\\sim5$) between the first and the third episode in order to explain the data is not a trivial assumption, especially since \\textit{all} other dSphs (including Carina) could be successfully modelled with a uniform SF efficiency. What can cause such a drastic change in the SF efficiency in Fornax and the interruption between SF episodes, and which galactic parameters control its SFH?  A hint for a possible explanation comes from the radial metallicity gradient in Fornax (e.g., \\citealt{Battaglia_06}). If the location of SF moves through the galaxy from outward in, and assuming that the SF efficiency depends on the radial position in the DM potential of the system, the net effect would be a time-dependence of the SF efficiency. In such a scenario, the star formation could be altered and stopped periodically by radiative heating of the ISM from SNe\\,II. Hydrodynamical simulations have shown, that SNe\\,II explosions can transfer sufficient thermal energy into the medium to produce local cavities in the density distribution and halt the SF process until the gas has cooled and fallen back to reignite the SF process (e.g., \\citealt{Ruiz_13}, \\citealt{dErcole_99}, see also \\citealt{Nichols_12} for an alternative explanation of periodic SF in dSphs).  Another evolutionary scenario is that Fornax experienced one, or several merger events with other gas-rich systems. For massive galaxies it has been shown that the SFR can be enhanced by a factor of up to 10, when they interact in close pairs or mergers (\\citealt{Scudder_12}). If this also holds for less massive systems, a merger scenario for Fornax could lead to a variation in the SF environment with time, and simultaneously explain the gap in stellar mass between Fornax and Sculptor, despite their apparently very similar early chemical enrichment history. In this case, the initial DM halo that defined the early chemical evolution of Fornax would have been less massive and might have formed initially only the metal-poor population of the MDF observed today. Through subsequent accretion events, the galaxy could have gained additional mass, and subsequently formed its enriched populations---with varying SF efficiencies.  In fact, there are several studies supporting a merger event for Fornax, from a number of observational aspects. \\citet{Coleman_04} identify shell-like overdensities in their photometry and \\citet{de_Boer_13} show that stars in these features are significantly younger than in its direct environment. MDFs constructed in \\citet{Battaglia_06} and \\citet{Coleman_08} show, that the distribution of stars is not homogeneous, but peak at several metallicities. \\citet{Battaglia_06} also find a distinct bimodality in the radial velocity distribution for the most metal-poor component, suggesting this subpopulation to be in an dynamically unrelaxed state. Note, that we find a similar distribution for this population in our kinematics (Hendricks et al. in prep). The distinct kinematic properties for different populations in Fornax found by \\citet{Amorisco_12} also support a merger scenario, and in fact these authors propose a ``bound-pair\" as a likely evolutionary scenario. In recent simulations dedicated to the complex kinematic and chemical structure observed in Fornax, \\citet{Yozin_12} are able to reproduce both the photometric overdensities as well as the peculiar velocity bimodality for MP stars in a scenario where Fornax has experienced a merger event between 3.5 and 2.1\\,Gyr ago. Note, however, that a bimodal velocity distribution in old stars, as observed in Fornax, may also be the result of stripped GC stars in a triaxial DM potential profile (\\citealt{Penarrubia_09}), whereas the fairly young shell structure ($\\sim 2$\\,Gyr) would require a significantly younger progenitor system.  Finally, \\citet{Kirby_11c} found that amongst several simple analytical models, their MDF of Fornax and other dSphs is best reproduced when the model allows for the infall of external (pristine) gas onto the galaxy during its evolution. The pure accretion of gas might have similar effects on the star forming environment in Fornax as mergers with a stellar component and therefore may also explain our findings.  The necessity of a bursty SFH with variations in the SF efficiency we propose here for the Fornax dSph arises from model predictions, and we do not know to what extent these models reflect the actual conditions in the galaxy. From observational side, the most promising way to test for the episodic nature of SF in Fornax---independent of its evolutionary interpretation---would be a deep photometric study of this galaxy, which is able to resolve the subgiant branch at $V\\approx23.5$\\,mag, and by that uncover possible distinct stellar population sequences. (e.g.,\\citealt{Smecker-Hane_96}, or \\citealt{Monelli_03} for Carina). Evidence for past mergers could be found in local overdensities of stars with the same velocity or [Fe/H] characteristic or galactic streams in the periphery of the galaxy. If the chemical enrichment efficiency inside Fornax changes with radius, we should expect to find different evolutionary sequences of, e.g., the $\\mathrm{\\alpha}$-elements from samples taken at different galactocentric distances. This, however, is a difficult endeavour since the fraction of metal-poor stars in the central region is extremely small.  On the other hand, the chemical evolution in a merger scenario is associated with an inflow of stars as well as pristine or enriched gas to an existing system. It is therefore beyond the capabilities of the models we use here, and new chemical and hydrodynamic models are needed in order to support or rule out different evolutionary pathways.  If Fornax actually grew from several smaller building blocks, it has to be asked whether this is an exception, or if it is a common evolutionary path for dwarf galaxies. In case Fornax instead did evolve in an isolated environment, and in fact did not gain a major part of its mass from external sources, the inevitable alternative is to introduce a new parameter to the chemical enrichment problem of this galaxy, such as a radial gradient of the SF and wind efficiency."
    },
    "1402/1402.3027_arXiv.txt": {
        "abstract": "Among 19 red-giant stars belonging to eclipsing binary systems that have been identified in \\textit{Kepler} data, 15 display solar-like oscillations. We study whether the absence of mode detection in the remaining 4 is an observational bias or  possibly  evidence of mode damping that originates from tidal interactions. A careful analysis of the corresponding \\textit{Kepler} light curves shows that modes with amplitudes that are usually observed in red giants would have been detected if they were present. We observe that mode depletion is strongly associated with short-period systems, in which stellar radii account for 16-24\\,\\% of the semi-major axis, and where red-giant surface activity is detected.  We suggest that when the rotational and orbital periods synchronize in close binaries, the red-giant component is spun up, so that a dynamo mechanism starts and generates a magnetic field, leading to observable stellar activity. Pressure modes would then be damped as acoustic waves dissipate in these fields. ",
        "introduction": "Binary systems hosting at least one star with detectable solar-like pulsations are becoming important astrophysical targets of particular interest because they provide a way to calibrate asteroseismology. As illustrated by the CoRoT and \\textit{Kepler} missions \\citep{Baglin_2009,Borucki_2010}, asteroseismology is an efficient method to measure masses and radii of large numbers of stars, which is of prime importance to test stellar evolution. However, a crucial test of direct comparison of  asteroseismic masses and radii of a large sample of stars with independent measurements of those quantites \\citep[e.g.][and refs.~therein]{Gaulme_2013b} has not yet been carried out.  Eclipsing binary systems (hereafter EBs) potentially permit such an exercise by allowing for accurate determination of masses and radii of both stars by combining photometric and radial-velocity measurements, provided that spectral lines are detectable for both components.  So far, all the stars known to both display acoustic modes and belong to EBs are red-giants (hereafter RGs), and all have been detected by the \\textit{Kepler} mission. The first detection was the 408-day period system KIC 8410637, whose mass was recently determined with radial-velocity measurements and agrees well with the asteroseismic result \\citep{Hekker_2010,Frandsen_2013}. Since then, \\citet{Gaulme_2013} reported a list of 17 \\textit{bona fide} new RG eclipsing-binary candidates (hereafter RGEBs), of which 12 displayed oscillations. Their orbital periods range from 15 to 987 days, their eccentricities from 0  to 0.7, and their companions are main-sequence stars (from spectral type L/M to F), except one that is  a second RG.  To be used as asteroseismic calibrators, the orbital periods of RGEBs must be measured with high precision, which implies that orbits be shorter than observation runs, i.e. $\\approx1500$~days for \\textit{Kepler}. More recently, \\citet{Beck_2013} reported the discovery of 18 new ``heartbeat'' stars, which are eccentric binaries undergoing dynamic tidal distortions \\citep{Welsh_2011}. Each of the 18 systems has a RG component displaying solar-like oscillations. The physical parameters of these systems were derived from modeling \\textit{Kepler}'s light curves and radial-velocity measurements.  Eclipsing binaries are also interesting for the physical processes resulting from tidal interactions, which may influence their evolution when the stars are close enough.  In this paper, we find evidence of a puzzling influence that binarity has on mode amplitudes and the surface activity of RGs. We consider all the known RGEBs, i.e. KIC 8410637, those from \\citet{Gaulme_2013}, and two new candidates. We aim to explore why, among the 19 systems, only 15 display solar-like oscillations. The peculiarity of the four cases with no RG pulsations is their rather short orbital periods ($\\leq 41$ days) and high quasi-periodic photometric variability (17 to 27\\,\\% in relative flux) indicative of large starspots. Taken together, this suggests that the absence of modes has four possible reasons: the modes are buried in the noise related to photometric variations, the mode frequencies are above the Nyquist frequency,  modes are damped because of activity due to strong binarity effects, or these RGs  are simply not pulsating. Damped modes would be consistent with \\citet{Chaplin_2011b}, who show that stellar activity inhibits the amplitude of solar-like oscillations.   We first present an update of the RGEB catalog by making use of about four years of \\textit{Kepler} data (quarters 0 to 16) to perform eclipse modeling and to estimate global asteroseismic parameters (Sec.~\\ref{sect_update}). Then we address the question of mode amplitudes and show evidence that mode damping must be occurring in conjunction with the presence of contrast features (spots, faculae) on the surface of the RGs that are tidally locked in phase (Sect. \\ref{sect_obs} \\& \\ref{sect_discussion}). ",
        "conclusions": "\\label{sect_discussion}  We have shown that the absence of detected solar-like oscillations in four of the 19 RGEBs is part of a bigger picture: the shorter the orbital period and the closer the stars are, the weaker the oscillation modes become, until complete mode depletion is reached. This is observed in the closest systems where rotational and orbital periods are almost synchronized and where strong surface activity is detected. Past works give context and help interpreting this original result.  \\citet{Chaplin_2011b} studied the impact of stellar surface activity on the detectability of solar-like oscillations for mostly main-sequence and sub-giant stars that were observed by \\textit{Kepler}. They measured the light-curve variability, in a similar manner as presented here, for about 500 solar-like stars that were observed at short cadence. They found evidence of a clear depletion of detected modes for stars with increased level of activity, attributing this to the presence of magnetic fields. They pointed out that magnetic structures in the solar photosphere are known to absorb acoustic oscillations and that strong magnetic fields can diminish the turbulent velocities in a convectively unstable layer, altering the mechanism for the  driving of acoustic modes. However, \\citet{Huber_2011} did not confirm this observation for the RGs of their sample.  For RGs, evidence of magnetic field and surface activity has been discussed  in the past, even though RGs were thought to be inactive due to decreased rotation during their late evolution. Single RG stars known as chromospherically active were identified by detection of Ca\\,{\\sc ii} H and K emission \\citep[e.g.][and refs. therein]{Dupree_Smith_1995,Auriere_2011,Konstantinova_2012}. The Zeeman Doppler imaging technique \\citep[][hereafter ZDI]{Donati_Brown_1997} allowed for mapping of the magnetic surface of such giants and revealed the presence of spots. In particular, the recent and preliminary work of \\citet{Konstantinova_2013} reports the detection of magnetic fields in 15 out of a sample of 29 G, K, and M single giants observed with ZDI. Active RGs likely constitute a minority of cases. The origin of their magnetic field and activity is thought to be generated either through dynamo mechanisms resulting from planet engulfment \\citep{Siess_Livio_1999}, angular dredge-up from the interior for the fastest rotators \\citep{Simon_Drake_1989}, or from remnant rotation for massive stars \\citep{Stepien_1993}. Active RGs were also identified in rather close detached  binary systems that usually rotate fast, with orbits and rotations in phase or close to a resonance, and where spots are detected with ZDI \\citep[][]{Ozdarcan_2010,Olah_2013}.   The case of the hierarchical triply eclipsing system HD~181068 \\citep[KIC 5952403,][]{Derekas_2011}, composed of a RG in a 45-day orbit with a pair of red dwarfs in a close 0.9-day orbit, displays common properties with our systems. \\citet{Fuller_2013} carefully analyzed 11 \\textit{Kepler} quarters of this system and could not detect any solar-like oscillations of the RG, even though its mass, radius, effective temperatures, and hence its $\\nu\\ind{max}$, are known from eclipse modeling and radial velocities. Their analysis indicates that the RG is tidally synchronized with the 45-day double dwarf orbit, and they suggest that a high magnetic activity resulting from such a rapid rotation suppresses the excitation of solar-like oscillations.  For the systems studied here, we  suggest  that binarity spins up RGs, with this phenomenon becoming stronger as systems are closer. This leads to the development of a dynamo mechanism, and thus the generation of  magnetic fields in the RGs that become visible at the surface. The resulting spots likely absorb part of the pressure mode energy making oscillations impossible to detect in the closest systems. Alternatively, we may think that the presence of spots shows that the convective energy is diverted into activity signal and not into global oscillations. This would mean that the properties of convection are considerably affected by binarity in the closest systems, and that oscillation excitation is reduced, or suppressed altogether. For the short-period systems KIC 8702921 and KIC 5308778, the orbits are almost circular, the rotational periods are harmonics of the orbital periods, and the orbital periods are typical of the other systems with no oscillations. We likely detect low-amplitude oscillations only because their companions are low-mass M-stars that generate rather weak tidal forces, as is also evident by the low amplitudes of their surface activity. Finally, this reasoning is consistent with the case of KIC 9246715, where the stellar masses are also a factor of two different and oscillations modes are detected only on the RG with larger mass.  We have discovered two new \\textit{bona fide} RGEBs and one heartbeat system candidate. KIC 8054233 is the longest-period RGEB in this work, and KIC 4663623 the system in which the temperature contrast  between the two components is the largest ($\\gtrsim 2$). We also updated the \\citet{Gaulme_2013} list by determining the orbital parameters of KIC 5640750, properly characterized the seismic parameters of KIC 5179609 whose $\\nu\\ind{max}$ is larger than the Nyquist frequency, detected oscillations in KIC 7377422, and removed a false-positive. Efforts in obtaining radial velocities of all systems and performing asteroseismic modeling on individual frequencies of RG oscillations are in progress in collaboration with the red-giant group of the Kepler asteroseismic science consortium.   \\appendix"
    },
    "1402/1402.3161_arXiv.txt": {
        "abstract": "We examine the exchange of angular momentum between the component spins and the orbit in semi-detached double white dwarf binaries undergoing mass transfer through direct impact of the transfer stream.  We approximate the stream as a series of discrete massive particles ejected in the ballistic limit at the inner Lagrangian point of the donor toward the accretor.  This work improves upon similar earlier studies in a number of ways.  First, we self-consistently calculate the total angular momentum of the orbit at all times.  This includes changes in the orbital angular momentum during the ballistic trajectory of the ejected mass, as well as changes during the ejection/accretion due to the radial component of the particle's velocity.  Second, we calculate the particle's ballistic trajectory for each system, which allows us to determine the precise position and velocity of the particle upon accretion.  We can then include specific information about the radius of the accretor as well as the angle of impact.  Finally, we ensure that the total angular momentum is conserved, which requires the donor star spin to vary self-consistently. With these improvements we calculate the angular momentum change of the orbit and each binary component across the entire parameter space of direct impact double white dwarf binary systems.  We find a significant decrease in the amount of angular momentum removed from the orbit during mass transfer, as well as cases where this process increases the angular momentum of the orbit at the expense of the spin angular momentum of the donor.  We conclude that, unlike earlier claims in the literature, mass transfer through direct impact need not destabilize the binary and that the quantity and sign of the orbital angular momentum transfer depends on the binary properties, particularly the masses of the double white dwarf binary component stars.  This stabilization may significantly impact the population synthesis calculations of the expected numbers of events/systems for which double white dwarfs may be a progenitor, e.g. Type Ia supernovae, Type .Ia supernovae, and AM CVn. ",
        "introduction": "\\label{sec-intro}  Double white dwarf (DWD) binaries provide an interesting key to understanding a number of different astrophysical problems.  Their birth properties provide insight into the evolution of their progenitors \\citep[e.g., ][]{2001A&A...375..890N}, as well as the dynamics of common envelope evolution \\citep[e.g., ][]{2012A&A...546A..70T}.  From a population standpoint, DWDs may make up a large fraction of the close binary stars in our galaxy, creating a confusion limited background for high frequency space-based gravitational wave detectors such as LISA \\citep{2001A&A...375..890N, 2010ApJ...717.1006R}.  As these objects evolve via gravitational radiation loses \\citep{1984ApJ...277..355W, 1984ApJS...54..335I} their shrinking orbit may eventually allow them to be detected as individual systems \\citep[e.g., ][]{1981ApJ...244..269N, 1987ApJ...323..129E, 1990ApJ...360...75H, 1988ApJ...334..947S}.  As the DWD orbit continues to shrink via gravitational radiation, the less massive component will inevitably fill its Roche lobe and the system will enter a semi-detached state.  It is in this interacting phase when sources and sinks of angular momentum other than gravitational radiation become vitally important in determining the eventual fate of the system.  \\citet{2004MNRAS.350..113M} \\citep[See also][in prep]{2007ApJ...655.1010G, 1999A&A...349L..17H, Kremer2013} analyze the stability of the mass transfer process for DWDs.  A number of possibilities have been discussed in the literature.  If the mass transfer process if unstable, the binary orbit will shrink until a merger occurs.  Such a merger may create such objects as R CrB or SdB stars \\citep[e.g., ][]{2002MNRAS.333..121S}.  If the system is sufficiently massive, a Type Ia supernova may result \\citep{1984ApJ...277..355W, 1984ApJS...54..335I, 2010ApJ...722L.157V, 2012A&A...546A..70T}.  Mergers of low-mass white dwarfs have also been proposed as the origin of Type .Ia supernovae \\citep{2010ApJ...715..767S, 2013arXiv1310.6359K}, where the properties of the binary immediately prior to merger can significantly affect the observed characteristics of such explosion \\citep{2010ApJ...722L.157V}. In the case where the mass transfer is stable, such systems may be identifiable as AM CVn \\citep{1996MNRAS.280.1035T, 2001A&A...368..939N}.  Three such stable double white dwarf AM CVn are thought to exist: HM Cnc, V407 Vul, and EX Cet \\citep[See][and references therein]{2010PASP..122.1133S}.  Each of these systems has an orbital period less than 12 minutes with no observable disk, making them excellent candidates for direct impact (DI) mass transfer. \\citet{2001A&A...368..939N} noted that, since the radius of the accretor in a DWD is large compared to the orbital separation, DI mass transfer is likely to occur.  Indeed, examining the distance of closest approach for the ballistic trajectories of \\citet{1975ApJ...198..383L}, inserting the Roche lobe radius of \\citet{1983ApJ...268..368E} and the mass-radius relation for zero-temperature white dwarfs of \\citet{1988ApJ...332..193V}, one finds \\citep[see Figure~1 of][Figure~\\ref{fig-di} of this paper]{2004MNRAS.350..113M} that the majority of the parameter space of DWDs will undergo DI mass transfer.  The evolution of the binary through this DI mass transfer phase crucially depends on the exchange of angular momentum between the binary components, as well as between the orbit and the component spins. Understanding this process is necessary to determine accurate rates for the creation of many of the astrophysical phenomena discussed previously.  \\citet{2004MNRAS.350..113M} finds that, while the stability of the system is largely dependent on the unknown tidal coupling between the accretor and the orbit, the majority of DI systems will be unstable.  \\citet{2007ApJ...655.1010G} notes that the stability of these systems may be improved if the rotation rate of the donor is allowed to vary.  However, both of these studies assume that (i) the orbital angular momentum of the particle at accretion is equal to the angular momentum of the circularization radius \\citep{1988ApJ...332..193V, 1975MNRAS.170..325F}, (ii) that all the angular momentum of the particle is deposited enitrely into the spin of the accretor, and (iii) that the angular momentum change during the particle's motion can be extracted from the semi-major axis of the binary orbit.  As we show here, these assumptions are not always justified.  Indeed, hydrodynamic studies of close white dwarf binaries suggest that these systems my be more stable than previously expected \\citep{2007ApJ...670.1314M, 2009JPhCS.172a2034D}.  In this paper, we follow the full ballistic trajectory of the ejected mass for each system, calculating the feedback onto the spin and orbital angular momenta of each star throughout the ejection, transit, and accretion process.  In section~\\ref{sec-ballist}, we describe our ballistic model and highlight the changes in the angular momenta of the donor and accretor.  In section~\\ref{sec-angmom} we examine the total changes in the angular momenta of the system for any DWD undergoing DI mass transfer and compare to previous results.  Finally, in section~\\ref{sec-disc}, we discuss the implications of these findings and propose a method to employ these results in future studies. ",
        "conclusions": "\\label{sec-disc}  We have shown that, if one takes into account the changes in both the spin and orbital angular momenta of the component stars during mass ejection/accretion, as well as numerically calculating the ballistic model for each mass transfer scenario, the resultant transfer of angular momentum can be significantly different than the result currently described in the literature.  In many cases the orbital angular momentum lost from the orbit can be significantly less than the standard assumption, making this process less destabilizing than expected.  This may allow for more DWD to survive the semi-detached state and become AM CVn, instead of merging to create Type Ia or Type .Ia supernovae.  This may have significant implications on population synthesis models of these objects.  In a few cases, we have shown that mass transfer may {\\it increase} the orbital angular momentum of the orbit, thereby providing a stabilizing effect on the orbit.  We caution again that the results presented here are for a single mass transfer event over a single binary orbit, and thus do not represent a steady state solution.  Nevertheless, any stabilizing effect increases the chances of a long-lived semi-detached DWD, lending credence to the creation of AM CVn through the DWD scenario.  In order to analyze changes in the long-term stability of these systems, we will incorporate our ballistic calculations into a long term numerical integration, taking into account tides and gravitational radiation.  By updating the ballistic trajectories and angular momentum changes at every timestep, we will be able to self-consistently calculate the long-term evolution of the system conserving the total angular momentum from ejection to accretion.  Results from these calculations will be presented in a forthcoming paper \\citep{Kremer2013}."
    },
    "1402/1402.6091_arXiv.txt": {
        "abstract": "{We examine whether the magnetorotational instability (MRI) can serve as a mechanism of angular momentum transport in circumplanetary disks. For the MRI to operate the ionization degree must be sufficiently high and the magnetic pressure must be sufficiently lower than the gas pressure. We calculate the spatial distribution of the ionization degree and search for the MRI-active region where the two criteria are met. We find that there can be thin active layers at the disk surface depending on the model parameters, however, we find hardly any region which can sustain well-developed MRI turbulence; when the magnetic field is enhanced by MRI turbulence at the disk surface layer, a magnetically dominated atmosphere encroaches on a lower altitude and a region of well-developed MRI turbulence becomes smaller. We conclude that if there are no angular momentum transfer mechanisms other than MRI in gravitationally stable circumplanetary disks, gas is likely to pile up until disks become gravitationally unstable, and massive disks may survive for a long time.} ",
        "introduction": "\\label{sec:1} Gas giant planets obtain mass from surrounding protoplanetary disks during their formation. When gas flows onto the planets, disks form around them as by-products. These gaseous disks are called circumplanetary disks. The evolution of circumplanetary disks is important not only for the formation of gas giants but also for the formation of satellites because regular satellites are thought to form in circumplanetary disks. There have been several theoretical studies of satellite formation in disks \\citep[e.g.,][]{lun82, can02, can06, mos03a, mos03b, can09, est09, sas10, ogi12}. \\citet{can02, can06} developed a so-called gas-starved disk model and successfully explained the total mass of the Galilean satellites and other satellite systems around giant planets in our solar system. In their scenario, gas and solids are continuously supplied to circumplanetary disks and satellites form and migrate onto the central planet repeatedly; the last generation of satellites remains when the mass inflow terminates. \\citet{mos03a, mos03b} developed a so-called solid enhanced minimum mass disk model and reproduced the Galilean satellites; the inner three satellites are formed in an inner massive disk and the outermost satellite is formed slowly in an outer extended low-density disk. A remaining problem is how to form circumplanetary disks accounting for their viscous evolution and mass infall from protoplanetary disks. Recently, a new idea for satellite formation has been advocated by \\cite{cri12}. Using their gas-free tidal-spreading particle disk model, the mass distributions of most regular satellite systems in our solar system can be well reproduced. However, it is difficult to formed Galilean satellites in this model, which would imply that the satellites still need a gaseous circumplanetary disk for their formation.  Hydrodynamic simulations of gas giant formation \\citep[e.g.,][]{lub99, tan02, ayl09a, ayl09b, mac06, mac08, mac10} have shown that circumplanetary disks form during the accretion phases of giant planets. Since most of these studies focused on gas giant formation, the detailed distribution of gas flow onto a circumplanetary disk was not well investigated. Recently, high-resolution 3D simulations by \\citet{tan12} have demonstrated that gas flows onto circumplanetary disks from high altitudes, not from the disk mid-plane. This was the first detailed analysis of the flux of gas infall from protoplanetary to circumplanetary disks. This picture is consistent with radiative hydrodynamical simulations by \\citet{kla06} and the very recent high-resolution global simulations by \\citet{gre13} and \\citet{ szu13}.  On the other hand, the driving mechanism of angular momentum transfer of gas in circumplanetary disks is not yet well understood. For accretion disks in general, the most promising mechanism is believed to be the magnetic turbulence driven by the magnetorotational instability (MRI). To be MRI-active, gas in a disk should be sufficiently ionized to couple with the magnetic field. The main ionization source in circumplanetary disks is galactic cosmic rays, and their attenuating length is about $100\\ {\\rm g\\ cm^{-2}}$. In some studies of MRI, a critical surface column density of $\\Sigma_{\\rm crit}\\sim100\\ {\\rm g\\ cm^{-2}}$ was adopted; a surface density below this value implies that the MRI can operate \\citep{gam96a}. However, if we take into account the chemical reactions of charged particles, such as recombination or capture by dust grains, the critical value can be far smaller \\citep*{san00, ilg06a, oku09, fuj11}. In the context of circumplanetary disks, \\citet{mar11a} and \\citet{lub12} pointed out the possibility of accretion outbursts induced by a combination of MRI and gravitational instability (GI).  In this work, we develop a model for circumplanetary disks by calculating the surface density with the mass infall rate obtained by \\citet{tan12}, and investigate whether the MRI is important in circumplanetary disks. Since \\citet{tan12} have found that the mass infall rate is proportional to the surface density of the parental protoplanetary disk, we can model the various evolution phases of the disk. We introduce a gas depletion factor and model the situation when a gas giant opens a gap, or gas in the protoplanetary disk is globally depleted. For the evaluation of MRI activity, we use the Elsasser number. The Elsasser number is proportional to the ionization degree, which we calculate for several conditions using the method developed by \\citet*{fuj11}. For ionization sources, we take into account galactic cosmic rays, X-rays from the host star, and the decay of short-lived radionuclides.  In Section \\ref{sec:2}, we explain how to estimate the border of MRI-active/inactive regions in gaseous disks such as circumplanetary disks or protoplanetary disks. Our models of circumplanetary disks are described in Section \\ref{sec:CPDmodel}. In Section \\ref{sec:MRIinCPD}, we investigate MRI activity in circumplanetary disks using our ionization degree calculation method described in Section \\ref{sec:2}. A discussion of our results is given in Section \\ref{sec:5}, and we summarize this paper in Section \\ref{sec:6}. ",
        "conclusions": "\\label{sec:5} We find that when accounting for cosmic rays, X-rays, and radionuclides, circumplanetary disks are not likely to sustain well-developed MRI turbulence. In contrast to protoplanetary disks, the volume ratio of MRI-active regions to dead zones for circumplanetary disks is very small. This is because the typical length scale of circumplanetary disks is smaller by several orders of magnitude than that of protoplanetary disks. This makes the timescale of magnetic diffusion smaller. According to Equation (\\ref{411}), even if $v_{{\\rm A}z}$ and $\\eta$ are the same, the Elsasser number tends to be smaller in circumplanetary disks. For example, the typical ionization degree at 5AU in a protoplanetary disk and that at 15$R_{\\rm J}$ in a circumplanetary disk are both $\\sim 10^{-10}$, but the Keplerian frequency of a protoplanetary disk at 5AU is $\\sim 10^{-8}$ s$^{-1}$ and that of a circumplanetary disk at 15$R_{\\rm J}$ is $\\sim 10^{-5}$ s$^{-1}$. Therefore, it is very difficult for MRI turbulence to be well developed in circumplanetary disks unless the surface density is very small, as we show in Fig. \\ref{fig:MRI}.  Recently, \\citet*{tur13} have investigated the possibility of the MRI in various models of circumplanetary disks from the literature. They choose the mid-plane value of plasma beta to be $10^3$. Their condition to have the MRI is that the magnetic pressure is smaller than the gas pressure, which is satisfied below 3.7 scale heights. They concluded that there are active layers at the disk surface. Our results are consistent with theirs when we choose the condition that $\\beta_{\\rm z}>8\\pi^2$ to sustain the MRI. For example, Figure 3 of \\citet{tur13} and Fig. \\ref{fig:f_dp=1} of this paper show similar models of surface density, and both have surface active layers. Note that even if the MRI can be sustained at the disk surface, it does not necessarily mean that there is well-developed turbulence. When MRI turbulence is well developed, the magnetically dominated atmosphere encroaches on a lower altitude and a region of well-developed turbulence becomes smaller \\citep{oku13a}. The main difference between \\citet{tur13} and our work is that we consider the criterion for turbulence to be well developed as well as that of just having MRI.  If there are no other mechanisms to give rise to viscosity and disks are not massive, the gas piles up in circumplanetary disks until the disks become massive enough to be gravitationally unstable. It is possible to promote gas accretion by GI, but it may not reduce the disk surface density much below the critical value for GI. Therefore the surface density is expected to remain large.  As we mentioned, if only GI can generate gas accretion, a massive and static disk will remain even after infall from the protoplanetary disk terminates. This suggests that the lifetime of circumplanetary disks may be longer than that of protoplanetary disks. Thus, perhaps we are more likely to be able to observe circumplanetary disks than previously thought \\cite[e.g.,][]{mam12}, and satellite formation may occur over a long timescale. On the other hand, it remains important to consider other mechanisms for angular momentum transport. A possible mechanism is spiral density waves caused by a non-axisymmetric potential \\citep[e.g.,][]{mac10, mar11b, riv12, szu13}, but this must be investigated in more detail.  If there is viscous heating in such a massive disk, it will easily heat up, and thermal ionization may drive the MRI \\citep{lub12}. Even if GI can drive turbulence, it does not necessarily mean the turbulence can generate heat in situ. The question of where the energy dissipates remains open \\citep{bal99, god01, mut10}. We should be careful in treating the heating by gravitational turbulence. Further study of energy dissipation by GI is necessary.  If thermal ionization triggers the MRI, disks would be less massive because of a high accretion rate. Suppose that if at each annulus of a disk, gravitational energy is converted into thermal energy and radiates as a black body from the disk surface, the effective temperature is \\begin{eqnarray} T_{\\rm eff} = \\left( \\frac{3GM_{\\rm p}\\dot M} {8\\pi\\sigma_{\\rm SB} r^3} \\right)^{1/4}, \\label{5-1} \\end{eqnarray} where $M_{\\rm p}$ is the planet mass, $\\dot M$ is the mass accretion rate, and $\\sigma_{\\rm SB}$ is the Stefan-Boltzmann constant. We can estimate the mid-plane temperature, $T_{\\rm c}$, from the approximation $T_{\\rm c} \\simeq \\tau^{1/4} T_{\\rm eff}\\ (\\tau \\gg 1)$, where $\\tau$ is optical depth, given by $\\tau\\sim\\kappa\\Sigma$ where $\\kappa$ is opacity. Here we use  $\\dot M = 3\\pi\\nu\\Sigma$. Then, the mid-plane temperature can be written as \\begin{eqnarray} T_{\\rm c} = 1.0\\times10^3\\left( \\frac{\\kappa} {5\\ {\\rm cm^2\\ g^{-1}}} \\right)^{1/5} \\left( \\frac{\\alpha}{10^{-2}} \\right)^{-1/5} \\left( \\frac{M_{\\rm p}}{M_{\\rm J}} \\right)^{3/10}\\nonumber\\\\ \\times \\left( \\frac{\\dot M} {3.1\\times10^{-7}M_{\\rm J}\\ {\\rm yr}^{-1}} \\right)^{2/5} \\left( \\frac{r}{10R_{\\rm J}} \\right)^{-9/10}\\ {\\rm K}, \\label{5-2} \\end{eqnarray} where ${M_{\\rm J}}$ is the Jupiter mass. When the mid-plane temperature exceeds about 1000K, gas in the disk will be sufficiently ionized to have the MRI. The inner disk may be hot enough to have thermal ionization especially during the early phases of gas giant formation; however, the outer disk seems to remain cool. More detailed calculations of the mid-plane temperature of circumplanetary disks have been done by \\citet{kei14}.  Another possibly important mechanism for satellite formation is the capture of planetesimals when they cross circumplanetary disks \\citep{fuj13} \\citep[see][for the gas-poor case]{est06}. Since our results suggest a large surface density, capture is expected to be effective. Non-axisymmetry in the density structure caused by these proto-satellites may play a role in angular momentum transport and it may be interesting to analyze that effect."
    },
    "1402/1402.6181.txt": {
        "abstract": "We present new results on the chemical composition of the Galactic ring nebula NGC~6888 surrounding the WN6(h) star WR136. The data are based on deep spectroscopical observations taken with the High Dispersion Spectrograph at the 8.2~\\unit{m} Subaru Telescope. The spectra cover the optical range from 3700 to 7400 \\AA. The effect of the CNO cycle is well identified in the abundances of He, N, and O, while elements not involved in the synthesis such as Ar, S, and Fe present values consistent with the solar vicinity and the ambient gas. The major achievement of this work is the first detection of the faint \\ion{C}{2} \\wav4267 recombination line in a Wolf-Rayet nebula. This allows to estimate the C abundance in NGC~6888 and therefore investigate for the first time the trace of the CNO cycle in a ring nebula around a Wolf-Rayet star. Although the detection of the \\ion{C}{2} line has a low signal-to-noise ratio, the C abundance seems to be higher than the predictions of recent stellar evolution models of massive stars. The Ne abundance also show a puzzling pattern with an abundance of about 0.5~\\unit{dex} lower than the solar vicinity, which may be related to the action of the NeNa cycle. Attending to the constraints imposed by the dynamical timescale and the He/H and N/O ratios of the nebula, the comparison with stellar evolution models indicates that the initial mass of the stellar progenitor of NGC~6888 is between 25~\\unit{M_\\odot} and 40~\\unit{M_\\odot}. ",
        "introduction": "\\label{intro} Massive stars ($\\geq20$~\\unit{M_\\odot}) play an important role in the chemical evolution and chemical enrichment of galaxies since they are major sources of metals in the Universe. They dominate the feedback to the local circumstellar environment during their relatively short lifetimes, between 3 and 10 million years \\citep{Ekstrometal12}. Through the emission of ionizing photons, the action of their strong stellar winds and their death as core-collapse supernova, massive stars interact with their surroundings injecting great amounts of radiation, mass and mechanical energy. An obvious manifestation of these interactions is the production of ring nebulae, i.e. filamentary shells that are found around evolved massive stars, especially Wolf-Rayet (WR) stars \\citep[e.g.][]{millerchu93, marstonetal94a, marstonetal94b, dopitaetal94, marston97, vanderhucht01}.  Actually, WR features revealing the interaction between massive stars and the interstellar medium (ISM) are often observed in the integrated spectrum of star-forming galaxies \\citep[e.g.][]{Schaereretal99, LopezSanchezEsteban09, LopezSanchezEsteban10a}. %%%%%%%%%%%%%%% \\begin{figure*} \\centering \\includegraphics[width=15.5cm]{f1.pdf} \\caption{The large extension of NGC~6888 is shown in the image obtained with the Wide Field Camera at the Isaac Newton Telescope (left, image release at http://www.ing.iac.es/PR/press/NGC6888.html). The image combines \\halpha\\ emission (red channel), 25\\% \\halpha\\ and 75\\% [\\ion{O}{3}] emissions (green channel), and [\\ion{O}{3}] emission (blue channel). The zoomed field to the right corresponds to an image observed with the Wide Field Planetary Camera 2 (WFPC2) on board the {\\it Hubble Space Telescope} by the 8568 program. The image represents our own combination of the filters F502N ([\\ion{O}{3}] \\wav5007, blue), F656N (\\halpha, green), and F658N ([\\ion{N}{2}] \\wav6583, red). The location of the slits observed with HDS are shown on the WFPC2 image. The slit covering the blue range is the one at the northwest and that corresponding to the red range is at the southeast. The slits were orientated with a position angle of 50$^\\circ$ and cover an area of 3\\arcsec$\\times$4\\farcs4.} \\label{f1} \\end{figure*} %%%%%%%%%%%%%%%  The true nature of ring nebulae was first pointed out by \\cite{johnsonhogg65} in the study of the Galactic objects NGC~6888 and NGC~2359. Today, their origin and evolution are understood in the framework of massive star evolution \\citep[e.g.][]{GarciaSeguraetal96b, GarciaSeguraetal96a, Freyeretal03, Freyeretal06, van-Marleetal05, van-Marleetal07}. It is well established that ring nebulae are interstellar bubbles of ionized gas that have swept up the surrounding medium during the mass-loss episodes of their central massive stars \\citep{weaveretal77}. The mass-loss episodes are simply the different phases of the stellar wind, which is especially powerful during the post-main sequence stages, playing an active role in the evolution of massive stars \\citep{Vinketal01, Vinkjsetal11}.  The interaction between the stellar wind and the circumstellar medium can produce ring nebulae with distinct characteristics. They will depend on the properties of the surroundings and of the stellar progenitor as well as its evolutive stage. Low-dense and hot interstellar bubbles are produced by the fast stellar winds of massive stars along the main sequence (MS) stage. They have been observed in the optical range \\citep[e.g.][]{marston97, stockbarlow10}, at $21$~\\unit{cm}  \\cite[e.g.][]{Cappaetal05} and at $24$~\\unit{\\mu m} \\cite[e.g.][]{Kraemeretal10}. These bubbles can be partially filled by the dense and slow stellar wind that operates during the phases of red supergiant (RSG) or luminous blue variable (LBV), depending on the initial mass of the progenitor. With the beginning of the WR phase, the material of the RSG wind is blown by the strong stellar wind and, finally, compressed into long filaments after the WR shell reaches the RSG shell. According to the calculations of \\cite{van-Marleetal05}, the remaining shell will dissipate over $\\sim80,000$ years after the beginning of the WR phase. In the literature, excellent works review in detail the evolutionary tracks of massive stars and the production of ring nebulae through the different stellar stages \\citep[e.g.][]{Arthur07, Freyeretal06, toalaarthur11}.  An interesting group of ring nebulae is that presenting chemical traces of processed stellar material;  the ejected and wind-blown types as defined in the classification of \\cite{chu81}, recently reviewed by \\cite{stockbarlow10}. These particular nebulae are thought to be blown by the stellar progenitors of long-duration $\\gamma$-ray bursts \\citep[e.g.][]{Vinketal11, Grafeneretal12}. They are powerful channels of information, directly connected to previous evolutionary phases. The enrichment process involves the transport of nucleosynthetic products from the stellar core to the outer layers, expelled in the post-MS stages. Thus, the study of their chemical abundances provides a unique opportunity to probe the stellar nucleosynthesis yield, an essential ingredient for our understanding of the chemical evolution of massive stars. From the analysis of He, N, O and Ne abundances, observational studies \\citep[e.g.][]{estebanetal90, estebanetal92, stocketal11} have shown that ring nebulae containing stellar ejecta have overabundances of He and N, and a deficiency of O. The results indicate that a substantial fraction of O has been transformed into N in the stellar interior via the activation of the ON cycle. The enrichment pattern is consistent with the stellar nucleosynthesis of the H-burning through the CNO cycle, the main nuclear reactions working in massive stars along the MS. However, the lack of reliable determinations of C abundances in any ring nebulae has hindered for decades the full knowledge of the CNO cycle trace.  NGC~6888 is a prototype example of Galactic ring nebulae: a wind-blown bubble presenting clear traces of stellar nucleosynthetic products \\citep[e.g.][]{kwitter81, estebanvilchez92}. A direct indication that NGC~6888 is a wind-blown bubble is the detection of X-ray emission in the bubble interior \\citep[e.g.][]{Bochkarev88, Wriggeetal94}, produced by the interaction between the WR wind and the RSG material that heats up the gas at $\\sim10^6$~\\unit{K} \\citep{weaveretal77}. Near the geometrical center of the nebula is located the stellar progenitor and ionizing source, the star HD~192163 \\citep[WR136;][]{vanderhucht01}, a WN6(h)-type star in the WR phase with a mass of about 15~\\unit{M_\\odot} \\citep{hamannetal06}. NGC~6888 is the best-studied nebula of its class thanks to its proximity \\citep[$1.45\\pm0.50$~\\unit{kpc};][]{van-Leeuwen07} and its large extension of about $8\\times5$~\\unit{pc^2} ($18'\\times12'$ on the sky). It resembles an almost complete ellipse, populated of filamentary structures of ionized gas. The mass of these structures amounts to 3.5--5~\\unit{M_\\odot} \\citep{wendkeretal75, kwitter81, marstonmeaburn88}, with a typical electron density of about 300 \\cmc\\ \\citep[e.g.][]{estebanvilchez92, fernandezmartinetal12}. The filaments mainly emit in [\\ion{N}{2}] and \\halpha, while they are surrounded by a highly-ionized, thin skin of material emitting in [\\ion{O}{3}] \\citep{mitra91, gruendletal00, mooreetal00}. The [\\ion{O}{3}] layer represents the boundary between the MS bubble and the WR shell; instead, the filaments are RSG wind material, possible mixed with the WR wind, compressed after the WR shell collides with the RSG shell \\citep[e.g.][]{kwitter81, estebanvilchez92, mooreetal00}. The gas kinematics have been studied in the past by several authors \\citep[e.g.][]{chu88, marstonmeaburn88}, showing that NGC~6888 expands on average at 80 \\kms, though the velocity field is complex with variations ranging from 55 to 110 \\kms\\ \\citep{lozinskaya70}. From these velocities, the dynamical age of the WR shell is estimated in $\\sim$30,000 years, varying between 20,000 and 40,000 years \\citep[see][]{Freyeretal06}.  In this paper, we review the CNO cycle trace in the chemical composition of the Galactic ring nebula NGC~6888, especially focusing on the determination of the total C abundance as derived from the faint \\ion{C}{2} \\wav4267 recombination line. In \\S\\ref{obsred} the observations are presented. A description of our spectrum as well as the methodology used in the line measurements and in the extinction correction can be found in \\S\\ref{med}. Physical conditions, chemical abundances and their calculations are found in \\S\\ref{results}. In \\S\\ref{chemo} the abundance pattern found in NGC~6888 is compared with those expected by the CNO cycle and the predictions by stellar evolution models. In \\S\\ref{discus}, the results are discussed in the light of previous ones for other similar objects and the future perspectives of exploring the C content in ring nebulae around massive stars. Finally, in \\S\\ref{conclu} we summarize our main conclusions.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conclu} We present results based on echelle spectroscopy of the Galactic ring nebula NGC~6888 taken with the High Dispersion Spectrograph at the 8.2~\\unit{m} Subaru Telescope. The nebula is associated to the central star WR136, a massive Wolf-Rayet (WR) star of the WN6(h) sequence. The spectra are very deep and cover the optical range from 3700 to 7400 \\AA. The high-spectral resolution of $R \\sim 12,000$ permits to distinguish three kinematical components at $-$60, $-$25, and +12 km$^{-1}$. The two blue-shifted components correspond to gaseous filaments or clumps in the approaching side of the nebula, and the red-shifted one to the ambient gas that surrounds the Cygnus OB1 association. We have derived the physical conditions, \\nel\\ and \\te, for each kinematical component. We obtain direct determinations of \\te([\\ion{O}{3}]) for the two blue-shifted components who belong to the shell of NGC~6888 itself, and an indirect estimation for the ambient gas. We determine ionic and total abundances of several elements from collisonally excited lines (CELs) as well as those of He and C from recombination lines (RLs). A comprehensive comparison of the gaseous abundances shows us that the effects of the CNO nucleosynthesis are well identified in He, N, and O, while elements such as Ar, S, and Fe --which are not involved in the cycle-- present abundances consistent with the solar values and the ambient gas. This is the first time that the \\ion{C}{2} \\wav4267 RL is detected in a ring nebula associated with a WR star, allowing to investigate the trace of the CNO cycle in a massive star ejecta. Although the detection of the \\ion{C}{2} line has a low signal-to-noise ratio, the C abundance seems to be higher than the predictions of recent stellar evolution models of massive stars, even including the effects of stellar rotation. Further investigations are needed making use of large aperture telescopes to obtain more reliable measurements for this faint emission line. The Ne abundance also presents a puzzling pattern, being lower than the model predictions and lower than the solar vicinity. These Ne deficiency may be also related to the Na overabundances observed in yellow supergiants. The action of the NeNa cycle occurring in a higher rate than the models assume may explain these results.  Attending to the constraints imposed by the dynamical timescale of the nebula and the observed He/H and N/O ratios, the comparison with the predictions of stellar evolution models indicates that the initial mass of the stellar progenitor of NGC~6888 should be between 25 and 40 M$_\\odot$. %We conclude that the stellar progenitor of the WR star WR136 probably had a mass between 32~\\unit{M_\\odot} with rotation and 40~\\unit{M_\\odot} without rotation. %Note that the future stellar study of WR136 (cite to Reyes-P\\'erez et al.) will constraint in a better way the properties of this star, which might be the next supernova type-I (?) in the next XX million years.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1402/1402.5571_arXiv.txt": {
        "abstract": "We report new observations of the prototype main-belt comet (active asteroid) 133P/Elst-Pizarro taken at high angular resolution using the Hubble Space Telescope.  The object has three main components; a) a point-like nucleus, b) a long, narrow antisolar dust tail and c) a short, sunward anti-tail.   There is no resolved coma. The nucleus has a mean absolute magnitude $H_V$ = 15.70$\\pm$0.10 and a lightcurve range $\\Delta V $ = 0.42 mag., the latter corresponding to projected dimensions 3.6$\\times$5.4 km (axis ratio 1.5:1), at the previously measured geometric albedo of 0.05$\\pm$0.02.  We explored a range of continuous and impulsive emission models to simultaneously fit the measured surface brightness profile,  width and position angle of the antisolar tail.  Preferred fits invoke protracted emission, over a period of 150 days or less, of dust grains following a differential power-law size distribution with index 3.25 $\\le q \\le$ 3.5 and having a wide range of sizes.   Ultra-low surface brightness dust projected in the sunward direction is a remnant from emission activity occurring in previous orbits, and consists of the largest ($\\ge$cm-sized) particles.   Ejection velocities of one micron-sized particles are comparable to the $\\sim$ 1.8 m s$^{-1}$ gravitational escape speed of the nucleus, while larger particles are released at speeds less than the gravitational escape velocity.  The observations are consistent with, but do not prove, a hybrid hypothesis in which  mass loss is driven by gas drag from the sublimation of near-surface water ice, but escape is aided by centripetal acceleration from the rotation of the elongated nucleus.  No plausible alternative hypothesis has been identified. ",
        "introduction": "In the traditional view, asteroids are rocky bodies formed inside the snow line of the Sun's proto-planetary disk while comets are ice-containing objects formed outside it. Recently, several hybrid objects having orbits interior to Jupiter's and the dynamical properties of asteroids, have been found to exhibit comet-like mass-loss (Hsieh and Jewitt 2006). Known as main belt comets (MBC) or, more generally, active asteroids (Jewitt 2012), these objects are candidate ice-bearing asteroids. Asteroid ice is of potentially far-reaching scientific interest.  Most importantly, the outer regions of the asteroid belt may have supplied some Terrestrial water and other volatiles, including the biogenic precursor molecules up to and including amino acids (Mottl et al.~2007). In the modern solar system, there is no known dynamical path linking stable orbits in the asteroid belt to the classical comet reservoirs in the Kuiper belt or Oort cloud. Instead, if ice exists in the asteroid main belt, it is likely to be primordial, trapped either at the formation epoch, or in some early, chaotic phase of solar system history (Levison et al.~2009).  About a dozen examples of active asteroids are currently known, but they do not all contain ice.  Data in-hand already clearly show the action of several different processes in causing mass-loss, ranging from asteroid-asteroid impact, to rotational break-up, to thermal fracture and more (Jewitt 2012).  Only for two objects, 133P/Elst-Pizarro (Hsieh et al.~2004, Hsieh and Jewitt 2006, Hsieh et al.~2010) and 238P/Read (Hsieh and Jewitt 2006, Hsieh et al.~2011), is the evidence for water sublimation reasonably strong.  The most compelling evidence is the repeated activity exhibited by both bodies. This is naturally explained by the seasonal sublimation of near-surface ice (including effects due to the orbital eccentricity as well as the obliquity of the nucleus, and possible self-shadowing by surface topography) but  is difficult or impossible to reconcile with the other mass-loss mechanisms so far envisioned.  No gas has yet been spectroscopically detected in any of the active asteroids, but the available limits to gas production are consistent with the very low ($<$1 kg s$^{-1}$) mass-loss rates inferred from dust (Jewitt 2012).  The evidence for ice sublimation thus remains indirect but no other plausible explanation for the observations, in particular for the seasonal recurrence of the activity, is known.  In this paper we describe the first high angular resolution, time-resolved images of 133P obtained using the Hubble Space Telescope. Our observations were triggered by the report of new activity on UT 2013 June 4, the fourth such episode of activity in 133P observed since 1996 (Hsieh et al.~2013).  We use the images to examine in detail the spatial and temporal properties of the mass-loss in the near-nucleus environment and to test the viability of the sublimating ice hypothesis. ",
        "conclusions": "\\label{discussion} \\subsection{Particle Properties: Mass,  Loss Rate and Lifetime} We estimate the dust mass in 133P as follows.  The mass of an optically thin collection of spheres of individual density $\\rho$ and mean radius $\\overline{a}$ is related to their total cross-section, $C_d$, by $M_d = 4/3\\rho \\overline{a} C_d$.  The effective mean radius of the particles contributing to the scattered light depends on the size distribution, $N(a)da$, the scattering area, $\\pi a^2$, and the time of residence of particles in the tail. The latter is size-dependent because radiation pressure imparts a velocity to dust particles that scales as $V$ = $V_1/a^{1/2}$, where $V_1$ is a constant. The amount of time that a particle of size $a$ spends in a given pixel of size $\\delta x$ is $t$ = $\\delta x/V$ = $(K a^{1/2})$, where $K$ is a constant.  Putting these together, we find that the effective mean particle radius weighted by the size distribution, the scattering cross-section and the time of residence is  \\begin{equation} \\overline{a} = \\frac{\\int_{a_{min}}^{a_{max}} a \\pi a^2 K a^{1/2} N(a) da}{ \\int_{a_{min}}^{a_{max}} \\pi a^2 K a^{1/2} N(a) da} \\end{equation}  \\noindent with $N(a)$ computed from Equation \\ref{Number}.  The continuous ejection models give 3.25 $\\le q \\le$ 3.5 (Table \\ref{parameter}). With these values, and assuming $a_{max} \\gg a_{min}$, we obtain  \\begin{equation} \\overline{a} \\sim \\frac{a_{max}}{5};~ (q = 3.25) \\label{3.25} \\end{equation}  \\begin{equation} \\overline{a} \\sim \\frac{a_{max}}{\\ln(a_{max}/a_{min})};~ (q = 3.5) \\label{3.5} \\end{equation}   \\noindent For example, in continuous emission starting at $t_0$ = 150 days, with 0.2 $\\le a \\le$ 10 mm (Table \\ref{parameter}), Equations (\\ref{3.25}) and (\\ref{3.5}) give $\\overline{a}$ = 2 to 3 mm.  With $\\rho$ = 1000 kg m$^{-3}$, we estimate the mass in the particle tail as $M_d \\sim$1.8$\\times$10$^7$ kg.   If this mass were released uniformly over  $\\sim$150 days, the implied mean mass-loss rate would be $dM_d/dt \\sim$ 1.4 kg s$^{-1}$, consistent with previous estimates of the mass-loss rates in 133P.  The corresponding values for the $t_0$ = 36 day solution are $\\overline{a}$ = 60 to 90 $\\mu$m, $M_d \\sim$5$\\times$10$^5$ kg and $dM_d/dt \\sim$ 0.2 kg s$^{-1}$.  We conclude that the continuous emission models bracket the mass-loss rates in the range 0.2 $\\le dM_d/dt \\le$ 1.4 kg s$^{-1}$.  These rates are 10 to 100 times larger than found by Hsieh et al.~(2004), largely because they used much smaller mean dust grain sizes than found here.   The mass-loss lifetime of the nucleus 133P can be estimated from  \\begin{equation} \\tau = \\frac{M_n}{f_o f_a (dM_d/dt)} \\label{tau} \\end{equation}  \\noindent in which $M_n$ is the mass of the nucleus, $f_o$ is the fraction of each orbit over which 133P ejects mass and $f_a$ is the duty cycle (ratio of active periods to elapsed time) for mass-loss between orbits.  The mass of a 2 km radius spherical nucleus of density $\\rho$ = 1300 kg m$^{-3}$ is $M_n \\sim$ 4$\\times$10$^{13}$ kg.  Observations show that $f_o \\sim$ 0.2 (Toth 2006) while $f_a$ is unknown.    Substituting into Equation (\\ref{tau}), and expressing $\\tau$ in years, we find 5$\\times$10$^6 \\le  \\tau f_a \\le$ 3$\\times$10$^7$ yr.  Strong upper limits to the duty cycle are obtained by setting $\\tau$ = 4.5$\\times$10$^9$ yr, the age of the solar system. Then, we must have $f_a <$ 0.001 for ice in 133P  to  survive for the age of the solar system.  This limit is consistent with most current (albeit widely scattered) observational limits from ensemble asteroid observations, according to which the instantaneous ratio of active to inactive asteroids is  1:6 to 1:60 (Hsieh 2009), $<$1:300 (Hsieh and Jewitt 2006), $<$1:400 (Sonnet et al.~2011), $<$1:25,000 (Gilbert and Wiegert 2010) and $<$1:30,000 (Waszczak et al.~2013).   It has been suggested that 133P might be a product of a recent ($\\lesssim$10 Myr) asteroid-asteroid collision (Nesvorny et al.~2009). If so, the preservation of ice would be trivial even for much larger duty cycles.   \\subsection{Ejection Mechanism}  The main observation suggesting that sublimation drives mass loss from 133P is the seasonal recurrence of activity.  Episodes of mass loss have repeated in the months near and immediately following perihelion in four consecutive orbits but mass loss was absent in between (Hsieh et al.~2004, Hsieh and Jewitt 2006, Hsieh et al.~2010, Hsieh et al.~2013).  Recurrence at a particular orbital phase strongly suggests a thermal trigger for the activity, and is reminiscent of the orbitally modulated mass loss observed in ``normal'' comets from the Kuiper belt.  None of the other mechanisms considered as drivers of mass loss from the active asteroids (including impact, rotational breakup, electrostatic ejection, thermal fracture and desiccation cracking) can produce seasonal activity on 133P in any but the most contrived way (Jewitt 2012).   Separately, the protracted nature of the emission in 133P is not readily explained by the other mechanisms but it is natural for sublimating ice.  Furthermore, we obtained good model fits to the data when assuming a drag-like size-velocity relation ($V \\propto a^{-1/2}$) but not when using a $V$ = constant velocity law.  The multi-parameter nature of the dust modeling problem means that we cannot formally exclude other size-velocity laws, but the success of the drag-like relation is at least consistent with a gas drag origin.  Lastly, Chesley et al.~(2010) reported a 3$\\sigma$ confidence detection of non-gravitational acceleration in 133P, most simply interpreted as a rocket effect from asymmetric sublimation.  However, the immediate problem for gas drag is that the solid particles are launched much more slowly than expected.  Micron-sized grains reach only $V_1 \\sim$ 1 m s$^{-1}$, as shown both by the order of magnitude approach in Section (\\ref{OOM}) and by the detailed models in Section (\\ref{details}) (see also Figure \\ref{width}).  This is quite different from active comets near the sun, in which micron-sized particles are dynamically well-coupled to the outflowing gas and attain terminal velocities comparable to the local gas velocity while even centimeter-sized grains exceed 1 m s$^{-1}$ (Harmon et al.~2004).    Whipple's (1951) model has become the first stop for estimating comet dust velocities (e.g.~Agarwal et al.~2007).  Applied to 133P, his model predicts the speed of micron-sized particles as $V_1 \\gtrsim$ 100 m s$^{-1}$ while millimeter grains should travel at 4.5 m s$^{-1}$, about two orders of magnitude faster than measured.  To examine the problem of the ejection speed more closely, we consider sublimation from an exposed patch of ice located at the subsolar point on 133P.  We solved the energy balance equation for a perfectly absorbing ice surface exposed at the subsolar point, with 133P located at 2.725 AU.  We included heating and cooling of the surface by radiation and cooling by latent heat taken up as the ice sublimates.  The resulting equilibrium sublimation mass flux is $F_s$ = 4$\\times$10$^{-5}$ kg m$^{-2}$ s$^{-1}$, which represents a maximum in the sense that the subsolar temperature is the highest possible temperature on the body.  The area of sublimating surface needed to supply  mass loss at rate $dM_d/dt$ [kg s$^{-1}$] is given by  \\begin{equation} \\pi r_s^2 = \\left(\\frac{1}{f_{dg} F_s}\\right) \\frac{dM_d}{dt} \\label{area} \\end{equation}  \\noindent where $r_s$ is the radius of a circle having area equal to the sublimating area and $f_{dg}$ is the ratio of dust to gas production rates.  In earlier work (Hsieh et al. 2004) we assumed $f_{dg}$ =1, as was traditional in cometary studies before long wavelength observations made possible the accurate measurement of dust masses in comets.  However, recent measurements of short-period comets have convincingly shown $f_{dg} >$ 1.  For example, infrared observations of comet 2P/Encke, whose orbit is closest to that of 133P among the short-period comets, give 10 $\\le f_{dg} \\le$ 30 (Reach et al.~2000).  Values $f_{dg} >1$ are physically possible because the ejected dust, although carrying more mass than the driving gas, travels much more slowly, allowing momentum to be conserved.  We conservatively take $f_{dg}$ = 10 and, with 0.2 $\\le dM_d/dt \\le$ 1.4 kg s$^{-1}$ use Equation (\\ref{area}) to obtain 500 $\\le \\pi r_s^2 \\le$ 3500 m$^2$, corresponding to a circular, sublimating patch on the surface with radius 13 $\\le r_s \\le$ 33 m.  Evidently, only a very tiny fraction of the nucleus surface  ($r_s^2/r_n^2 \\sim$ (1 to 5)$\\times$10$^{-5}$) of exposed ice is needed to supply the dust mass loss in 133P.   The tiny size of the active area affects the dust speed, as we now show.  Outflow of the gas into the surrounding half-space vacuum will produce a  wind that exerts a drag force on entrained dust particles.  If outgassing occurred from a point source, the gas density would fall in proportion to $d^{-2}$, where $d$ is the distance from the source.  If outgassing occurred from an infinite plane, the gas density would vary as $d^0$, for the same reason that the surface brightness of an extended light source is independent of the distance from which it is viewed.  The intermediate case, in which the source is neither a point nor an infinite plane, requires a full gas dynamic treatment to determine the iso-density surfaces in the expanding gas, which lies beyond the aims of the present work.  Instead, we have made a ``small source approximation'' (SSA) model as described in the Appendix.  Its solution is plotted in Figure (\\ref{V_vs_a}), where it is compared with the Whipple model.    For both the SSA and the Whipple models, we plot two curves.  The upper curves (solid lines) show the terminal dust speeds in the absence of nucleus gravity while the lower curves (dashed lines) show the full solutions including gravity.   The SSA velocity law has a form similar to Whipple's formula for the speed of an escaping particle, but differs in that the source size, $r_s$, now plays a role in the terminal velocity, with smaller $r_s$ corresponding to lower ejection speeds.  Physically, the difference between our result and Whipple's is the length scale over which gas drag accelerates entrained dust particles.  Whipple assumed that sublimation was uniform from the entire sun-facing hemisphere and so, in his model, the length scale is the radius of the nucleus, $r_n$.  In our model, the length scale is the (much smaller) size of the source, $r_s$.    In simple terms, the difference between the models is the difference between a handgun and a rifle; the latter fires a faster bullet than the former because the longer barrel gives a greater acceleration length over which momentum from the explosive gases is transferred to the projectile.  While the SSA model gives lower speeds than the Whipple model, the predicted grain speeds are still about an order of magnitude larger than those measured in 133P (Figure \\ref{V_vs_a}).  We see several possibilities for explaining this difference. One possibility is that $F_s$ is  over-estimated because the ice is not located exactly at the subsolar point.  Ice might also be protected beneath a thin refractory mantle where it would receive only a fraction of the full solar insolation, reducing $F_s$.  Of course, a severe reduction in the sublimation flux would further reduce the critical grain size, $a_c$, magnifying the puzzle of how millimeter and centimeter sized grains can be ejected.  Another possibility is that sublimation proceeds, not from a single patch of  radius $r_s$, but from a large number of  sources  individually small compared to $r_s$, but having a combined area $\\pi r_s^2$, such that $V_T$ from Equation (\\ref{VT}) is further reduced.   This might be a natural result of progressive mantle growth on an aging ice surface, as larger blocks left behind on the surface restrict sublimation to a network of small, unblocked areas.  By Equation (\\ref{VT}), an order of magnitude reduction in dust speed to match the measured values would be produced by sublimation from individual sources of scale $r_s \\sim$ 0.2 m.  Yet another possibility is that ejected dust particles leave the surface containing some fraction of attached ice, the anisotropic sublimation of which would exert a reaction (``rocket'') force, propelling the grains laterally out of the gas flow.  Small, spherical dust grains should be essentially isothermal, giving no net force, but an irregular distribution of ice within a porous, aspherical, rotating grain  will in general lead to a net sublimation force even from an isothermal particle.  A very modest ice fraction in each grain would be sufficient to propel the particles laterally over distances $\\gtrsim r_s$, especially in 133P where $r_s \\sim$ 20 m is so small.  The process would be less important on more active (generally better studied) comets because of their larger source regions.  Unfortunately, the data offer no way to determine whether these effects, working either separately or collectively, occur in 133P.  It is highly improbable that gas drag acting alone could launch particles from the nucleus into the dust tail with terminal velocities (i.e.~velocities at distances $r \\gg r_n$) much smaller than the gravitational escape speed  from the nucleus.  In order to  climb out of the gravitational potential well of the nucleus while retaining a terminal velocity $V_T(a)$, dust grains must be launched from the surface at speed $V$ = $(V_T(a)^2 + V_e^2)^{1/2}$, where $V_e \\sim$ 1.8 m s$^{-1}$ is the gravitational escape speed from 133P.  For example, millimeter grains by Equation (\\ref{vperp}) have $V_T(a) \\sim$ 6 cm s$^{-1}$, from which the above relation gives a launch velocity $V  \\sim$ 1.801 m s$^{-1}$ (i.e.~only 0.05\\% larger than $V_e$).  Gas drag alone is unable to provide such fine-tuning of the ejection velocity over a wide range of particle sizes.  The solution to this problem may be that gravity is nearly or completely negated by centripetal acceleration on parts of 133P, as a result of  the elongated shape and 3.5 hr rotation period.   We envision very weak gas flow driven by sublimation near the  tips of the elongated nucleus, perhaps in the form of a fluidized bed on which gravity is reduced nearly to zero.  In this scenario, the particles would accelerate by gas drag according to a $V \\propto a^{-1/2}$ velocity-size relation, consistent with the data.  Their escape would be unimpeded by gravity so that even the largest, slowest particles could escape to be swept up by radiation pressure.  Forces other than gas drag might eject particles under these circumstances, but only sublimation can account for the observed seasonal modulation of the activity on 133P, as originally proposed (Hsieh et al.~2004). In this regard, 133P is a C-type object (Hsieh et al.~2004) and, as noted above, must have density $\\rho_n \\gtrsim$ 1300 kg m$^{-3}$ (Equation \\ref{density}) if regolith material is to be gravitationally bound at its rotational extremities.  We note that the average density of C-type asteroids is 1300$\\pm$600 kg m$^{-3}$  (Cary 2012), consistent with $\\rho_n$.    A hybrid solution in which gas drag provides the force to expel particles with rotation reducing the potential barrier is thus plausible, although not proved.     \\clearpage"
    },
    "1402/1402.0448_arXiv.txt": {
        "abstract": "We discuss a nonlocal modification of gravity obtained adding a term $m^2 R\\,\\Box^{-2}R$ to the Einstein-Hilbert action. We find that the  mass parameter $m$ only affects the non-radiative sector of the theory, while the graviton remains massless, there is no propagating ghost-like degree of freedom, no vDVZ discontinuity, and  no Vainshtein radius below which the theory becomes strongly coupled. For $m={\\cal O}(H_0)$ the theory therefore recovers all successes of GR at solar system and lab scales, and only deviates from it at cosmological scales.  We examine the cosmological consequences of the model and we find that it  automatically  generates a dynamical dark energy and a self-accelerating evolution. After fixing our only free parameter $m$ so to reproduce the observed value of the dark energy density today, we get a pure prediction for the dark energy equation of state, $w_ {\\rm DE}\\simeq-1.14$. This value is consistent with the existing data,  and could also resolve the possible tension between the Planck data and local measurements of the Hubble parameter. ",
        "introduction": "The experimental observation of the accelerated expansion of the Universe \\cite{Riess:1998cb,Perlmutter:1998np} has stimulated an intense search for modifications of General Relativity (GR) at cosmological scales. The construction of a consistent infrared deformation of GR turns out however to be extremely challenging. A natural way to proceed is to introduce a mass scale $m$ of the order of the present value of the Hubble parameter $H_0$. However such attempts must face a number of difficulties, related to the possible appearance of new ghost-like degrees of freedom (or of ghostlike excitations over non-trivial backgrounds), the appearance of classical or quantum strong coupling regimes, potential problems with causality, while at present it is also unclear whether acceptable cosmological solutions emerge (see \\cite{Hinterbichler:2011tt,deRham:2014zqa} for  recent  reviews).  In a recent series of papers \\cite{Jaccard:2013gla,Maggiore:2013mea,Foffa:2013vma,Foffa:2013sma,MMAK} our group has proposed an approach in which a mass parameter enters the theory as the coefficient of a suitable nonlocal term. At the level of the general idea,  our approach was inspired by the observation that nonlocal operators provide a way of writing  a mass term, both in massive electrodynamics and in linearized massive gravity, without breaking the gauge invariance of the massless theory~\\cite{Dvali:2006su} and they can also play an important cosmological role through the degravitation mechanism  \\cite{ArkaniHamed:2002fu}. In practice, this general idea can be implemented in different ways. The one closest to the original degravitation idea involves the addition of a term  $m^2(\\iBox\\Gmn)^{\\rm T}$ to the Einstein equations \\cite{Jaccard:2013gla}. The  superscript T  denotes the extraction of the transverse part, which is necessary for consistency with energy-momentum conservation (see also \\cite{Porrati:2002cp} for related ideas). It was then realized in \\cite{Maggiore:2013mea,Foffa:2013vma,Modesto:2013jea} that such tensor nonlocalities generate instabilities in the cosmological evolution. Similar conclusions were obtained in \\cite{Ferreira:2013tqn} studying a nonlocal model with a term of the form $\\Rmn\\iBox\\RMN$ in the action. In refs.~\\cite{Maggiore:2013mea,Foffa:2013sma,Foffa:2013vma,MMAK} we have then turned our attention to a model in which a term $m^2(\\gmn\\iBox R)^ {\\rm T} $ is added to the Einstein equations, and we found that it passes a number of tests of theoretical consistency, and  has an interesting cosmological phenomenology. In this paper we turn our attention to a related model, in which again the $\\iBox$ operator acts on the Ricci scalar, but which is defined by the nonlocal action \\be\\label{S1} S_{\\rm NL}=\\frac{1}{16\\pi G}\\int d^{d+1}x \\sqrt{-g}\\, \\[R-\\frac{d-1}{4d} m^2R\\frac{1}{\\Box^2} R\\]\\, , \\ee where  $d$ is the number of spatial dimensions and the factor $(d-1)/4d$ is a convenient normalization of the mass parameter $m$. We will see that this model is quite interesting, both at the theoretical level and for its cosmological consequences. Non-local cosmological models of different type, not involving a mass scale, have also been  studied recently \\cite{Deser:2007jk,Nojiri:2007uq,Jhingan:2008ym,Koivisto:2008xfa,Koivisto:2008dh,Capozziello:2008gu,Elizalde:2011su,Zhang:2011uv,Elizalde:2012ja,Park:2012cp,Bamba:2012ky,Deser:2013uya,Ferreira:2013tqn,Dodelson:2013sma,Woodard:2014iga,Barvinsky:2003kg,Barvinsky:2011hd,Barvinsky:2011rk}. ",
        "conclusions": ""
    },
    "1402/1402.1602_arXiv.txt": {
        "abstract": "In this paper, we report for the first time the detailed temperature evolution process of the magnetic flux rope in a failed solar eruption. Occurred on January 05, 2013, the flux rope was impulsively accelerated to a speed of $\\sim$400 km/s in the first minute, then decelerated and came to a complete stop in two minutes. The failed eruption resulted in a large-size high-lying ($\\sim$100 Mm above the surface) high-temperature ``fire ball\" sitting in the corona for more than two hours. The time evolution of the thermal structure of the flux rope was revealed through the differential emission measure analysis technique, which produced temperature maps using observations of the Atmospheric Imaging Assembly on board \\textit{Solar Dynamic Observatory}. The average temperature of the flux rope steadily increased from $\\sim$5 MK to $\\sim$10 MK during the first nine minutes of the evolution, which was much longer than the rise time (about three minutes) of the associated soft X-ray flare. We suggest that the flux rope be heated by the energy release of the continuing magnetic reconnection, different from the heating of the low-lying flare loops, which is mainly produced by the chromospheric plasma evaporation. The loop arcade overlying the flux rope was pushed up by $\\sim$10 Mm during the attempted eruption. The pattern of the velocity variation of the loop arcade strongly suggests that the failure of the eruption be caused by the strapping effect of the overlying loop arcade. ",
        "introduction": "Magnetic flux ropes play an important role in solar eruptions manifested as coronal mass ejections (CMEs) and/or flares. Recently, a new line of observational structures, namely EUV hot-blobs and/or channels, have been proposed to be the most direct manifestation of flux ropes (Cheng et al. 2011; Zhang et al. 2012; Cheng et al. 2013; Patsourakos et al. 2013) with the Atmospheric Image Assembly (AIA; Lemen et al. 2012) on board the \\textit{Solar Dynamics Observatory (SDO)} (Pesnell et al. 2012). In particular, Zhang et al. (2012) revealed the flux rope as a conspicuous hot channel structure prior to and during a solar eruption, which initially appeared as a writhed sigmoid with a temperature as high as $\\sim$10 MK, then continuously transformed itself toward a semi-circular shape and acted as the essential driver of the resulting CME.  While previous studies have been focusing on the morphological and kinematic evolution of magnetic flux ropes, little is known about their  thermal evolution. The scenario of heating solar flare loops is well accepted in classical flare models (Carmichael 1964; Sturrock 1966; Hirayama 1974; Kopp \\& Pneuman 1976), in which flare loops are mainly heated by the thermal plasma from the chromospheric evaporation (Doschek et al. 1980; Feldman et al. 1980). Now the question is whether flux ropes, seen as expanding hot channels in a much higher corona, are heated by the same process. In this letter, we intend to use differential-emission-measure (DEM) based temperature analysis method on an event of failed flux rope eruption to address this issue. Recently, DEM analyses have been applied to diagnose the physical properties of CMEs (Zhukov \\& Auch\\`{e}re 2004; Landi et al. 2010; Tian et al. 2012; Cheng et al. 2012; Tripathi et al. 2013), the quiet Sun (V\\'{a}squez et al. 2010) and the post-flare loop systems (Reeves \\& Weber 2009; Warren et al. 2013). The multi-passband broad-temperature capability of AIA makes it ideal for constructing DEM models. We find that this method is particularly useful in studying the failed solar eruption, in which the thermal structure is conspicuous and long lasting in AIA field of view (FOV).  There have been several studies on failed eruptions, but the physical cause of the failure remains elusive. It is found that long duration flares tend to be more eruptive (Kahler et al. 1989), but many  other factors could be invovled, such as the interaction between filaments and their associated magnetic environment (Ji et al. 2003; Williams et al. 2005; Gibson \\& Fan 2006; Jiang et al. 2009; Kuridze et al. 2013), the degree of the helical twist in the filament (Rust \\& LaBonte 2005), the overlying arcade field (Wang \\& Zhang 2007), and the position of the reconnection site (Gilbert et al. 2007). Numerical simulations showed that the kink instability could trigger a failed filament eruption if the overlying magnetic field decreases slowly with height (e.g., T\\\" or\\\" ok \\& Kliem, 2005).  For the failed eruption in this paper, we can clearly observe and track the flux rope structure, as a hot EUV blob, before, during and after the eruption. Using advanced DEM-based temperature map method, we reveal, for the first time, the detailed thermal evolution of the flux rope, and conclude that the thermal structure is caused by the direct heating from the energy release of the magnetic reconnection in the flux rope. Further, we find convincing evidences that the failure of the eruption is caused by the strapping effect of the overlying magnetic loop arcade. In Section 2, we present the observations and results, which are followed by a summary and discussions in Section 3. ",
        "conclusions": "A failed flux rope eruption associated with an M1.7 class flare was observed by the AIA at the northeast limb of the Sun on 2013 January 5, which provided us an unprecedented opportunity for studying the detailed temperature evolution of the flux rope. We find that the flux rope existed in the corona with a high temperature ($\\sim$5 MK) before its eruption. The temperature evolution of the flux rope can be divided into two phases: impulsive-heating phase and post-impulsive-heating phase, which correspond well with the rise phase and decay phase of the associated flare. On the contrary, the temperature of the low-lying flare loops only increased during the flare rise phase, and quickly decreased during the decay phase. Furthermore, there was a gap of low temperature between the flux rope and the flare loops. All of these observations indicate that the flux rope should be heated by a different process from  that of the flare loops. The heating process of flare loops have been well understood and widely accepted as summarized below. Energetic electrons and ions produced by flare magnetic reconnection precipitate from the coronal acceleration site and lose their energy in the dense underlying chromosphere via Coulomb collisions. The temperature in the chromosphere increases and the resulting pressure exceeds the ambient chromospheric pressure, which leads the heated plasma to expand along the magnetic field, forming flare loops (Doschek et al. 1980; Feldman et al. 1980). While for the flux rope, we suggest that it should be heated directly by the thermal energy generated at the reconnection site through the thermal conduction. It is unlikely that chromospheric evaporation plays a strong role in heating the flux rope.  The velocity analysis of the flux rope and the overlying loop arcades strongly suggests that the failure of the eruption be caused by the strapping effect of the overlying magnetic loop arcade. The strapping effect, or the tension force of the line-tying field, is the determining factor that contained the flux rope from further eruption."
    },
    "1402/1402.6352_arXiv.txt": {
        "abstract": "We extend the statistical analysis of Lissauer et al.~(2012, ApJ 750, 112), which demonstrates that the overwhelming majority of \\ik candidate multiple transiting systems (multis) represent true transiting planets, and  develop therefrom a procedure to validate large numbers of planet candidates in multis as {\\it bona fide} exoplanets.  We show that this statistical framework correctly estimates the abundance of false positives already identified around \\ik targets with multiple sets of transit-like signatures based on their abundance around targets with single sets of transit-like signatures. We estimate the number of multis that represent split systems of one or more planets orbiting each component of a binary star system. We use the high reliability rate for multis to validate more than one dozen particularly interesting multi-planet systems % herein.  Hundreds of additional multi-planet systems % are validated in a companion paper by Rowe et al. (2014, ApJ ).  We note that few very short period ($P < 1.6$ days) planets orbit within multiple transiting planet systems and discuss possible reasons for their absence. There  also  appears to be a shortage of planets with periods exceeding a few months in multis. ",
        "introduction": "More than 40$\\%$ of the \\ik planet candidates announced by \\cite{Borucki:2011}, \\cite{Batalha:2013} and \\cite{Burke:2013} are associated with targets that have more than one candidate planet.  % Accounting for candidates on each one's individual merit, without considering whether the same target has any additional candidates, \\cite{Morton:2011} estimated the fidelity of \\ikt's planet candidates (fraction of the candidates expected to be actual planets) to be above $90\\%$. Incorporating the corrections presented in \\cite{Santerne:2013}, \\cite{Fressin:2013} derive a slightly larger fraction of false positives (FPs), although they classify planets orbiting stars other than the \\ik target to be FPs. If one doesn't consider such planets to be FPs, then their estimate of the number of FPs is reduced by more than a factor of two, leading to a planet fidelity rate of $\\sim 95\\%$. In this work, we do not classify planets orbiting a star that is bound to the \\ik target to be FPs, and among the candidates that we validate as planets, we expect very few, if any, planets orbiting  stars unrelated to the \\ik target. The galactic latitude distribution of \\ik targets with both single and multiple planet candidates tracks that of \\ik targets, but is quite different from that of false positives produced by chance-alignment blends with faint eclipsing binary stars (Appendix A), increasing our confidence in the \\ik planet candidate sample.  \\ik has found far more multiple planet candidate systems (multis) than would be the case if candidates were randomly distributed among target stars \\citep{Lissauer:2011b, Latham:2011}. \\cite{Lissauer:2012}, henceforth Paper I, presented a statistical analysis that combined the large numbers of multis observed by \\ik  (as listed in \\citealt{Borucki:2011}) together with the assumption that false positives are nearly randomly distributed among \\ik targets to demonstrate that  the fidelity of \\ik multiple planet candidates is far higher than that for singles.  We expand upon the statistical  analysis of Paper I herein and show that it correctly accounts for the number of FPs that have been identified in multis.  The high reliability of the remaining sample forms the foundation of a framework that we develop to validate a large number of \\ikt 's multiple planet candidates as  true planets. We present the techniques that we use to search for companions of planet-hosting stars, and then validate and discuss  multi-planet systems of high multiplicity and/or special dynamical interest.  A companion paper by \\cite{Rowe:2014}, henceforth Paper III, % describes the light curve studies and astrometric measurements of transit locations using pixel-level \\ik data to reveal the position of the center of light (centroid) of the \\ik target star both during transits and outside of transit, along with the classification of planet-hosting stars. Paper III validates hundreds of new multiple planet systems.%  In Section 2, we refine the statistical analysis presented within Paper I. We discuss the ensemble of candidate planetary systems considered for validation, which is based upon multis identified using the first two years of \\ik data, in Section 3. We combine the equations of Section 2 with the candidate pool presented in Section 3 to yield numerical estimates of the number of FPs present in multis in Section 4. Section 4.2 describes tests of the statistical framework for assessing the expected number of FPs in multis using the distributions of \\emph{identified} FPs and remaining planet candidates. In Section 5, we consider and classify the various sources that can produce multiple, periodic, transit-like signatures in the light curves of \\ik target stars, with the principal sources being transiting planets and eclipsing binary stars. We provide estimates of the fraction of multis that represent blended planetary systems (including planets orbiting each of the components of a binary star system)  in Section 6. The criteria that we use to validate \\ik planet candidates as {\\it bona fide} planets are presented in Section 7. In Section 8, we present  our search for stellar multiplicity in these systems. The three planets within binary star system Kepler-132 = KOI-284, with one of the stars possessing two transiting planets and the other star being orbited by a single transiting planet, as well as the five planets on S-type (circumstellar) orbits in the Kepler-296 = KOI-1422 binary system, are validated and discussed in Section 9. Section 10 summarizes the properties of other \\ik systems with five or more transiting planet candidates, beginning with the seven-planet Kepler-90 = KOI-351 system in Subsection 10.1.  The  Kepler-223 = KOI-730 planetary system with four planets very close to a ``chain'' of mean motion resonances is presented in Section 11.   % We conclude the main text by summarizing our principal results in Section 12.  Appendix A compares the distributions in galactic latitude of \\ik targets, planet candidates and false positives. Appendix B discusses the differing period distributions of various classes of transit-like signatures that have been identified in \\ik data and the implications of these differences for planet validation and planetary system architecture. Computing false positive rates by comparing overall \\ik planet candidate abundances to abundances of eclipsing binary stars severely underestimates the reliability of multis but overestimates the reliability of single planet candidates. The ramifications of the large abundance and high reliability of multis on estimates of the reliability of \\ik singles  are addressed in Appendix C.  Appendix D presents the  dynamical simulations that we performed to assess stability of systems with tightly-packed planet candidates as well as those to determine which of the candidate planets that are large enough to be brown dwarfs can be demonstrated to be planets by limits on transit timing variations (TTVs) induced in neighboring planets. ",
        "conclusions": "A sizable majority of the planet candidates in \\ik multis are validated as planets herein and in Paper III.  Statistical validation of \\ik multis leads to a high level of confidence of planet status for each of the validated objects.   The vast majority of the objects passing our tests are planets orbiting the \\ik target; combined with a much smaller group of planets orbiting bound stellar companions to the \\ik target stars, we expect that  more than $99\\%$ of the objects that we validate  herein and in Paper III are planets that travel around the Galaxy together with the star that we associate with them.  Table \\ref{tab:fpest} provides estimates of the number of FPs of various types expected among the 1033 members of the multiple planet candidate population that satisfy the six criteria listed in Section 3 and thus are considered for validation in this study. Under the pessimistic assumption that 10\\% of the \\ik single planet candidates are FPs, the total expected number of FPs is a mere 2.1. The most likely type of FP among the multiple planet systems that we have validated is a nominal two-planet system that actually contains 1 FP and 1 planet; a very small number of nominal two-planet systems may actually be 2 FPs.  The most likely pollutants of systems of higher multiplicity are 1 FP plus two or more planets; it is unlikely that any of the  systems of 3 or more validated planets contains more than one FP, although that possibility cannot be strictly excluded.  Provided that one does not consider planets transiting a stellar companion to the \\ik target as FPs, the \\cite{Fressin:2013} and \\cite{Santerne:2013} studies suggest an FP rate for singles of only 7\\%, even after the upward adjustment of their estimates that neglect planetary multiplicity to apply to the subpopulation of singles as described in Appendix C.   Using an FP rate for singles of 7\\% rather than 10\\% reduces our estimate of the number of FPs in multis by more than 30\\%.  These estimates are for the entire population of currently viable (candidate) multis that satisfy the six criteria listed in Section 3; the subset of multis validated herein and in Paper III is smaller and so the equivalent estimates are also smaller. The bottom line is that unless orbital planes of planetary systems are highly correlated with those of close binaries that they are bound to within triple (and higher multiplicity) star systems, $\\gtrsim 99.8\\%$ of the objects that we validate should be true transiting exoplanets.  Although our work provides a catalog consisting of a large number of new, highly-reliable, planetary systems, it is far from complete, and it is biased in favor of common characteristics that makes systems amenable to our methods of statistical validation.  \\emph{If a planet candidate has not been validated, that does not imply that it isn't a real planet; indeed, we expect that over 90\\% of \\ik multis not validated by our studies (nor identified to date as FPs) are nonetheless actual planets and planetary systems.}  Thus, while our sample is useful for those who want to study characteristics of individual systems without worry of contamination by false positives, the entire set of \\ik candidate multi-planet systems is less biased and still highly reliable, making it superior for use in statistical studies of the distribution of planetary system characteristics."
    },
    "1402/1402.6955_arXiv.txt": {
        "abstract": "In this work the processes of absorption charge-exchange and photo-association in He+H$^{+}$ collisions together with the process of ion HeH$^{+}$ photo-dissociation are considered as factors of influence on the opacity of the atmospheres of helium-rich white dwarfs in the far UV and EUV region. It is shown that they should be taken into account even in the cases of the atmospheres of white dwarfs with H:He =$10^{-5}$. Than, it is established that in the cases of white dwarfs with H:He $\\gtrsim 10^{-4}$, particulary when H:He $\\approx 10^{-3}$, these processes have to be included \\emph{ab initio} in the corresponding models of their atmospheres, since in the far UV and EUV region they become dominant with respect to the known symmetric ion-atom absorption processes. ",
        "introduction": "It has been shown recently in \\citet{mih13}, that in order to consider the contribution of the absorbtion processes connected with binary ion-atom systems to the opacity of the solar photosphere it is not enough to take into account only the processes of absorbtion charge exchange in (H + H$^{+}$)-collisions and the molecular ion H$_{2}^{+}$ photo-dissociation. These processes were studied in \\citet{mih86} and \\citet{mih93a, mih94b, mih07a}, and are already included in one of the solar photosphere models \\citep{fon09}. It has been established that in the very important far UV and EUV spectral regions they have to be considered together with the processes of the absorbtion charge exchange and photo-association in non-symmetric (H + X$^{+}$)-collisions and molecular ion H$X^{+}$ photo-dissociation, where $X$ is one of metal atoms. Namely, it has been proved that only in such case the total efficiency of the ion-atom absorbtion processes in the mentioned spectral regions approaches the efficiency of the relevant concurrent processes in the whole solar photosphere.  These results suggest that it is useful to consider again the situation of ion-atom absorbtion processes in the atmospheres of helium-rich white dwarfs. Let us remind that in the previous papers \\citep{mih92b, mih94a, sta94, mih95, ign09}, dedicated to certain DB white dwarf atmospheres, the processes of molecular ion He$_{2}^{+}$ photo-dissociation were studied: \\begin{equation} \\label{eq:sim1} \\varepsilon_{\\lambda} + He_{2}^{+} \\longrightarrow He + He^{+}, \\end{equation} and absorbtion charge exchange in (He + He$^{+}$)-collisions: \\begin{equation} \\label{eq:sim2} \\varepsilon_{\\lambda} + He^{+} + He \\longrightarrow He + He^{+} \\end{equation} where $\\varepsilon_{\\lambda}$ is the energy of a photon with wavelength $\\lambda$, He = He$(1s^{2})$, He$^{+}$ = He$^{+}(1s)$ and He$_{2}^{+}$ = He$_{2}^{+}(1^{2}\\Sigma_{u}^{+})$. The significance of these symmetric ion-atom absorption processes for the atmospheres of the considered DB white dwarfs was established in \\citet{mih94a, mih95} and \\citet{ign09} by a direct comparison of their efficiencies with the main concurrent process of inverse \"bremssthrallung\" in (free electron + He)-collisions, i.e. \\begin{equation} \\label{eq:e-He} \\varepsilon_{\\lambda} + e + He \\longrightarrow e' + He, \\end{equation} where $e$ and $e'$ denote a free electron in the initial and final energetic states respectively. For that purpose the data from the corresponding DB white-dwarf atmosphere models \\citep{koe80} have been used. It was established that the processes (\\ref{eq:sim1}) and (\\ref{eq:sim2}) significantly influence the opacity of the considered DB white dwarf atmospheres, with an effective temperature $T_{eff} \\ge 12000$ K, which fully justifies their inclusion in one of the models of such atmospheres \\citep{ber95}. However, the same comparison demonstrated also that the dominant role in those atmospheres generally still belongs to the concurrent absorbtion process (\\ref{eq:e-He}), while the processes (\\ref{eq:sim1}) and (\\ref{eq:sim2}) can be treated as dominant (with respect to this concurrent process) only in some layers of that atmospheres, and only within the part 50 nm$ < \\lambda < $ 250 nm of the far UV and EUV region. In Fig. \\ref{fig:plank}, where Plank's curves for $T_{eff}=12000$ K and $14000$ K are shown, this part is denoted by $\"I\"$. Its boundary (from the short-wavelength side) is determined by the value of wavelength $\\lambda_{He} \\approx 50.44$ nm, which corresponds to the threshold of the atom He photo-ionization. Hence it follows that in the case of helium-rich white dwarf atmospheres it would certainly be useful to include into consideration some new ion-atom absorbtion processes, which is principally allowed in accordance with the composition of such atmospheres \\citep{bue70}. \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth, height=0.75\\columnwidth]{plank.pdf} \\caption{Plank curve for $T_{eff.}=12000$ K and $T_{eff.}=14000$ K. \"I\" and \"II\" denote the regions $50.44$ nm $\\le \\lambda \\lesssim 250$ nm and $50.44$ nm $\\le \\lambda \\lesssim 120$ nm respectively. } \\label{fig:plank} \\end{center} \\end{figure}  Let us note in this context that in the case of white dwarf atmospheres with dominant helium component, among all possible symmetric ion-atom absorbtion processes which are allowed by their composition \\citep{bue70}, only the processes (\\ref{eq:sim1}) and (\\ref{eq:sim2}) have to be taken into account. This means that in this case we can find new relevant absorbtion processes only among the processes connected with non-symmetric ion-atom systems, particulary such systems which could provide efficiency in the same part $\"I\"$ of the far UV and EUV region. Here we will examine the significance of non-symmetric ion-atom absorption processes with participation of the hydrogen component. We mean the processes of molecular ion HeH$^{+}$ photo-dissociation \\begin{equation} \\label{eq:nonsim1} \\varepsilon_{\\lambda} + HeH^{+} \\longrightarrow He^{+} + H \\end{equation} and the processes of absorption charge exchange and photo-association in the (He + H$^{+}$)-collisions, namely \\begin{equation} \\label{eq:nonsim2} \\varepsilon_{\\lambda} + He + H^{+} \\longrightarrow He^{+} + H, \\end{equation} \\begin{equation} \\label{eq:nonsim3} \\varepsilon_{\\lambda} + He + H^{+} \\longrightarrow (HeH^{+})^{*} , \\end{equation} where H = H$(1s)$, He= He$(1s^{2})$, He$^{+}$ = He$^{+}(1s)$, and HeH$^{+}$ and (HeH$^{+}$)$^{*}$ denote the molecular ion in the ground and the first excited electronic states which are adiabatically correlated with the states of the systems He + H$^{+}$ and He$^{+}$ + H respectively at the infinite internuclear distance. Already in \\citet{mih13} it was noted that these processes, whose significance was practically neglected for the solar photosphere, could be rather important in the case of helium-rich white dwarf atmospheres. This assumption was worthy of attention, particulary due to the fact that characteristics of the considered non-symmetric molecular ions (see Fig.\\ref{fig:mol.ion}) provide manifestations of the processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}) just in the part $\"I\"$ of the far UV and EUV region.  In connection with this fact, let us remind that the part $\"I\"$ is rather important for such values of $T_{eff}$. Namely, let $\\lambda_{max}$ be the positions of the maxima of the spectral intensities characterizing the electromagnetic (EM) emission of the considered atmospheres, which are determined from the well-known Wien's law: $\\lambda_{max} \\cdot T_{eff} = 2.898 \\cdot 10^{6}$ nm $\\cdot$ K. We then have it that for the considered DB white dwarfs $\\lambda_{max} < 250$ nm, so that the mentioned maxima lie just within the region $\"I\"$.  It was just because of the above mentioned that this investigation was undertaken. The main aim was to study when the non-symmetric absorbtion processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}) can significantly influence the opacity of helium-rich white dwarf atmospheres in the part $\"I\"$ of the UV and VUV spectral region, and to show that the processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}) deserve to be included \\emph{ab initio} in the corresponding white dwarf atmosphere models. Therefore the relevant spectral characteristics of the processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}) are determined here for the atmospheres of different helium-rich white dwarfs with $T_{eff}=12000$ K and $14000$ K, log $g=8$ and $7$, and for the values of the ratio H:He from $10^{-5}$ to $10^{-3}$. The necessary expressions for these spectral characteristics are given in Section 2. Then, with their help, in the Section 3 the values are calculated of the parameters characterizing the relative efficiency of the non-symmetric processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}) with respect to the efficiency of all ion-atom processes, as well as with respect to the electron-atom process (\\ref{eq:e-He}), which comprise the main direct results of this work.  Let us note that already in \\citep{ign09}, beside the electron-atom processes (\\ref{eq:e-He}) the process of photoionization of hydrogen atoms, was also considered, namely \\begin{equation} \\label{eq:phi} \\varepsilon_{\\lambda} + H \\longrightarrow e + H^{+}, \\end{equation} where H=H$(1s)$. It was treated as a concurrent process, potentially necessary in the region $\\lambda < \\lambda_{H}$ where $\\lambda_{H}\\approx 911$ {\\AA} corresponds to the threshold of atom H photoionization. However, in \\citep{ber13} it was brought to attention that the importance of this absorption channel was significantly underestimated. That is why in this work the process (\\ref{eq:phi}) was again included into the consideration and carefully examined.  One can see that in this work we take into account only such non-symmetric ion-atom absorption processes where, apart from the dominant helium component of the considered atmospheres, only the hydrogen component participates, although they contain also a lot of metal components \\citep{bue70}. This is due to the fact that the existing atmosphere models do not provide the necessary data (about the relevant metals' abundances) which would be needed for the present calculations. However, we consider that a demonstration of the fact that for the considered atmospheres the processes (\\ref{eq:nonsim1}) - (\\ref{eq:nonsim3}), where one of their \"minor\" components participates, are rather significant, is a sufficient reason for treating the non-symmetric ion-atom absorption processes in general as potentially significant for those atmospheres. ",
        "conclusions": "From the presented material it follows that the considered non-symmetric ion-atom absorption processes have to be treated as one of the important channels of influence on the opacity of the atmospheres of helium-rich white dwarfs in the far UV and EUV region. So, it has been shown that even in the case of DB white dwarfs with H:He = $10^{-5}$ such processes should be included in the models of their atmospheres. However, the main result of this research is the establishment of the fact that in the cases of helium-rich white dwarfs with H:He $> 10^{-5}$, and particularly with H:He =$10^{-4}$, these non-symmetric ion-atom absorption processes have to be included \\emph{ab initio} in the models of the corresponding atmospheres, since in the greater part of the considered far UV and EUV region they could be completely dominant with respect to the referent electron-atom and symmetric ion-atom absorbtion processes. \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth, height=0.75\\columnwidth]{kappa-3-4-5.pdf} \\caption{The behavior of the total ion-atom spectral absorption coefficient $\\kappa_{ia}(\\lambda;\\log \\tau)$, see Eqs. (\\ref{eq:kapansim}), (\\ref{eq:kapasim}) and (\\ref{eq:kapaia}), in the cases of the atmospheres of helium rich white dwarfs with $T_{eff} = 12000$ K and $\\log g=8$ for: H:He = $10^{-5}$, H:He = $10^{-4}$ and H:He =$10^{-3}$.} \\label{fig:kappa_ia5} \\end{center} \\end{figure}  Besides, attention has been payed again in this paper to the role of the hydrogen component in the atmospheres of helium-rich white dwarfs. Namely, it has been shown that in all the considered cases (H:He = $10^{-5}, 10^{-4}$ and $10^{-3}$) the hydrogen photo-ionization processes (\\ref{eq:phi}) yield a dominant contribution to the opacity of the corresponding atmospheres in the region $\\lambda_{He} < \\lambda < \\lambda_{H}$.  As a task for further investigations in this area the study can be mentioned of the atmospheres of helium-rich white dwarfs with smaller effective temperatures ($\\approx 7000$ K) where the significance of the hydrogen component can be greater than in the described cases. Also, inclusion into consideration would be useful of the ion-atom non-symmetric absorption processes with participation of some metal components of the considered atmospheres."
    },
    "1402/1402.4482_arXiv.txt": {
        "abstract": "We study outflows driven by Active Galactic Nuclei (AGNs) using high-resolution simulations of idealized $z\\sim2$ isolated disk galaxies.  Episodic accretion events lead to outflows with velocities $>1000$\\kms and mass outflow rates of about the star formation rate (several tens of \\MsunPerYear).  Outflowing winds escape perpendicular to the disk with wide opening angles, and are typically asymmetric (i.e. unipolar) because dense gas above or below the AGN in the resolved disk inhibits outflow.  Owing to rapid variability in the accretion rates, outflowing gas may be detectable even when the AGN is effectively ``off.''  The highest velocity outflows are sometimes, but not always, concentrated within $2-3$~kpc of the galactic center during the peak accretion.  With our purely thermal AGN feedback model -- standard in previous literature -- the outflowing material is mostly hot ($\\gtrsim 10^6$~K) and diffuse ($n_H \\lesssim 10^{-2}$~cm$^{-3}$), but includes a cold component entrained in the hot wind.  Despite the powerful bursts and high outflow rates, AGN feedback has little effect on the dense gas in the galaxy disk.  Thus AGN-driven outflows in our simulations do not cause rapid quenching of star-formation, although they may remove significant amounts of gas over long ($\\gtrsim$Gyr) timescales. ",
        "introduction": "Many models of the co-evolution of galaxies and supermassive black holes (BHs) invoke active galactic nuclei (AGN)-driven winds to regulate black hole growth and impact the host galaxy \\citep[e.g.][]{ciotti97, haehnelt98, silk98, fabian99, ciotti01, wyithe03, begelman05, dimatteo05, sazonov05, springel05_mergers_ellipticals, king05, hopkins06_unified, robertson06a, somerville08, ciotti10, king10a, king10b, dubois10,novak11, teyssier11, debuhr12, dubois12_dual, martizzi12, dubois13b, dubois13}. Such outflows may help establish correlations between the black hole and galaxy properties \\citep[e.g.][]{kormendy95, magorrian98, richstone98, ferrarese00, gebhardt00, mclure02, tremaine02, marconi03, treu04, peng06, graham07, treu07, jahnke09, bennert11, canalizo12, graham13, kormendy13}.  Powerful AGN winds may help expel sufficient quantities of gas from their host galaxies to regulate their star-formation, and in this way suppress the stellar mass growth of high-mass galaxies.  Observations have begun to probe galaxy-scale outflows thought to be driven by AGN feedback \\citep{tremonti07, feruglio10, alatalo11, lehnert11, nesvadba11, rupke11, aalto12, cicone12, cano-diaz12, harrison12, maiolino12, cimatti13, forster-schreiber13, gonzalez-alfonso13, rupke13, spoon13, veilleux13, cicone14}.  These outflows are detected in both ionized and molecular gas, and from the local Universe up to $z\\sim2$.  The highest velocities are often found along lines-of-sight near the galactic nucleus, and these velocities ($\\gtrsim10^3$\\kms) seem to be higher than those caused by stellar feedback \\citep[e.g.][]{rupke05, martin05, rupke11}.  These high velocities and central locations are frequently used to argue that the outflows are powered by AGN rather than supernovae or young stars.  With some assumptions, mass outflow rates are estimated at $\\sim 100$\\MsunPerYear or sometimes even higher, or a few times the star-formation rate \\citep[SFR; e.g.][]{cicone14, forster-schreiber13}.  Some authors suggest that outflows thought to be AGN-driven are quenching star-formation in their host galaxies by expelling the gas, yet most AGN host galaxies are still forming stars.  Most observed AGNs occur in undisturbed star-forming disk galaxies up to redshift $z\\sim2$, and AGNs are quite common in these galaxies \\citep{grogin05, pierce07, gabor09, georgakakis09, silverman09, cisternas11, kocevski12, mullaney12, rosario12, schawinski12, simmons12, juneau13}, although especially luminous or obscured AGNs may be more likely to be hosted by merging galaxies \\citep{bennert08, cales11, bessiere12, ramos-almeida12, treister12, juneau13, hopkins13}.  It remains unclear whether moderate-luminosity AGNs in disk galaxies can drive outflows or have a major impact on the evolution of their host galaxies.  Observed outflows attributed to AGN in fact sometimes occur in galaxies with little or no indications of current AGN \\citep{tremonti07, forster-schreiber13}.  In \\citet{gabor13} we presented simulations of gas-rich, isolated disk galaxies whose supermassive black holes accrete at a high rate. Accretion is highly variable, with stochastic $\\sim10$~Myr episodes of Eddington-limited accretion.  In the present paper, we show that these high-accretion episodes generate powerful AGN outbursts, driving hot, high-velocity outflows from the galaxy.  The outflow rates peak near the SFR, but the outflows have little overall effect on the gas in the galactic disk.  We summarize our simulation methods in \\S\\ref{sec.sims}, describe our results in \\S\\ref{sec.results}, and discuss caveats and implications in \\S\\ref{sec.discussion}.  We conclude with \\S\\ref{sec.conclusion}. ",
        "conclusions": "\\label{sec.conclusion} We have studied AGN-driven outflows using a small suite of simulations of gas-rich, high-redshift isolated disk galaxies.  The main results from our simulations are:  \\begin{itemize} \\item AGN outbursts occur in short, episodic bursts lasting $\\sim5-10$~Myr. \\item AGN feedback drives high-velocity outflows with mass outflow rates peaking near the SFR (a few 10's of \\MsunPerYear).  The time-averaged outflow rate can be $\\sim1/3$ of the SFR. \\item Dense gas in the galactic disk directs AGN outflows out of the plane of the disk.  The outflows propagate into the circum-galactic region as semi-spherical blast waves. \\item AGN outflows are frequently asymmetric, expanding only above or below the gas disk, not both. Dense gas clouds vertically offset from the central engine can block the outflow in some directions. \\item AGN outflows are predominantly hot and diffuse ($T\\sim10^7$~K and $n_H\\sim10^{-3}$~cm$^{-3}$), but some colder, denser clouds ($T\\sim10^2-10^4$~K and $n_H \\sim 10^{-1}$~cm$^{-3}$) are swept up in the hot outflows. These detailed properties of outflows, along with outflow rates, are somewhat sensitive to the AGN feedback model.  Since the high-accretion events that drive outflows have injection temperatures above $10^7$~K in any case, our results are \\emph{not} particularly sensitive to the to minimum injection temperature of our feedback model. \\item During peak BH accretion, the highest outflow velocities are sometimes, but not always, visible along lines-of-sight within $2-3$~kpc of the galactic center.  In some cases high velocities are visible far from the galactic center due to the transverse expansion of outflowing shells. \\item AGN feedback has only a weak effect on the detailed dynamics of the galaxy's gas.  The mass of dense gas in the disk is almost identical for simulations with and without AGN feedback. \\end{itemize}  Beyond these main points, we have speculated on how these results could be model-dependent, and we have discussed their relation to observations and the popular idea of AGN-driven quenching.  We advocate a picture where typical AGN-driven outflows contribute to the overall galactic outflows, but do not affect the dense gas in the disk or the SFR.  AGNs fueled by disk instabilities exist in a quasi-steady state where powerful outflow bursts balance nuclear inflows over long timescales, without quenching star-formation across the galaxy disk."
    },
    "1402/1402.2571_arXiv.txt": {
        "abstract": "The Crab nebula pulsar was observed in 2009 January and December with a novel very fast optical photon counter, Iqueye, mounted at the ESO 3.5 m New Technology Telescope. Thanks to the exquisite quality of the Iqueye data, we computed accurate phase coherent timing solutions for the two observing runs and over the entire year 2009. Our statistical uncertainty on the determination of the phase of the main pulse and the rotational period of the pulsar for short (a few days) time intervals are $\\approx 1 \\, \\mu$s and $\\sim$0.5 ps, respectively. Comparison with the Jodrell Bank radio ephemerides shows that the optical pulse leads the radio one by $\\sim 240 \\, \\mu$s in January and $\\sim 160 \\, \\mu$s in December, in agreement with a number of other measurements performed after 1996. A third-order polynomial fit adequately describes the spin-down for the 2009 January plus December optical observations. The phase noise is consistent with being Gaussian distributed with a dispersion $\\sigma$ of $\\approx 15 \\, \\mu$s in most observations, in agreement with theoretical expectations for photon noise-induced phase variability. ",
        "introduction": "\\label{sec1}  Only a handful of pulsars show optical emission at detectable levels and, among them, only five show pulsations in the optical band: the Crab (PSR B0531+21; PSR J0534+2200) and Vela (PSR B0833-45) pulsars, PSR B0540-69, PSR B0656+14, and Geminga (PSR B0630+17) (see e.g. \\citealt{mignani11}). The first source in which such pulsations were detected is the Crab nebula pulsar \\citep{cocke+69,lynds+69}. As such, the study of these pulsars, and in particular of the Crab which is the brightest among them ($V\\sim 16.6$ mag), is of noticeable importance to understand the optical emission mechanism and to investigate, by comparison with the radio and other wavebands, the geometry of the emission regions and of the magnetic field.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Log of the 2009 observations of the Crab nebula pulsar, taken with Iqueye mounted at the ESO 3.5 m NTT telescope in Chile. Times refer to the barycentre of the Solar system (TEMPO1).} \\label{tab1} \\begin{tabular}{@{}llrrrrlrlr@{}} \\hline &  Observation ID  &  Start time (UTC)  & Start time (MJD) & Duration (s) \\\\ \\hline 1 & 20090115-030226UTC$_{-}$crab2 & January 15, 03:11:14.9 & 54846.132811 &  255.2  \\\\ 2 & 20090115-031933UTC$_{-}$crab3 & January 15, 03:28:07.8 & 54846.144535 &  198.2  \\\\ 3 & 20090115-032345UTC$_{-}$crab4 & January 15, 03:32:18.8 & 54846.147440 &  897.8  \\\\ 4 & 20090116-014452UTC$_{-}$crab & January 16, 01:53:40.9 & 54847.078946 &  156.9  \\\\ 5 & 20090116-014802UTC$_{-}$crab & January 16, 01:56:48.7 & 54847.081119 &   98.2  \\\\ 6 & 20090116-015011UTC$_{-}$crab & January 16, 01:59:10.9 & 54847.082765 &  233.7  \\\\ 7 & 20090116-015448UTC$_{-}$crab2 & January 16, 02:03:37.6 & 54847.085852 &  1997.6 \\\\ 8 & 20090117-012137UTC$_{-}$crab & January 17, 01:30:18.3 & 54848.062712 &  3597.5  \\\\ 9 & 20090118-014557UTC$_{-}$crab & January 18, 01:55:02.2 & 54849.079887 &  727.1   \\\\ 10 & 20090118-020018UTC$_{-}$crab & January 18, 02:09:22.3 & 54849.089842 & 5997.2   \\\\ 11 & 20090119-004736UTC$_{-}$crab & January 19, 00:56:09.6 & 54850.039000 &  3598.1\t\\\\ 12 & 20090120-012735UTC$_{-}$crab & January 20, 01:36:16.4 & 54851.066857 &  3712.7  \\\\ 13 & 20091213-034209UTC$_{-}$crab-NTT & December 13, 03:51:47.5 & 55178.160966 &  178.3  \\\\ 14 & 20091213-035043UTC$_{-}$crab-NTT & December 13, 04:00:33.5 & 55178.167054 & 1798.2  \\\\ 15 & 20091213-042211UTC$_{-}$crab-NTT & December 13, 04:31:47.5 & 55178.188744 &    6.1  \\\\ 16 & 20091214-023608UTC$_{-}$crab-NTT & December 14, 02:46:08.7 & 55179.115379 & 7198.2  \\\\ 17 & 20091214-043954UTC$_{-}$crab-NTT & December 14, 04:49:32.8 & 55179.201074 & 7198.3  \\\\ 18 & 20091215-031945UTC$_{-}$crab-NTT & December 15, 03:29:34.9 & 55180.145543 & 7198.6  \\\\ 19 & 20091216-021802UTC$_{-}$crab-NTT & December 16, 02:27:33.9 & 55181.102476 & 3598.5  \\\\ 20 & 20091216-031834UTC$_{-}$crab-NTT & December 16, 03:28:05.9 & 55181.144513 &  898.3  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  The optical light curve of the Crab pulsar is well known. It is characterized by a double-peak folded profile, separated in phase by $\\sim 0.4$ (or $\\sim 140^0$). It is very stable \\citep{zampieri+11}, although marginal evidence for a secular decrease in luminosity \\citep{nasuti+96} and weak variations in the pulse shape \\citep{karpov+07} have been reported.  Absolute optical timing of the Crab pulsar light curve has been addressed in some past investigations. For the other optical pulsars, with the sole exception of PSR B0540-69 \\citep{middleditch+87,gradari+11}, the weakness of the targets prevented this type of studies till now. After the early attempts of \\citet{nelson+70}, \\citet{papaliolios70} and \\citet{horowitz+71}, long-term (2 through 10 years) timing solutions for the Crab pulsar in the optical band were obtained by \\citet{boynton+72}, \\citet{groth75a,groth75c}, and \\citet{lohsen81}, who found that the pulsar spin-down has a complex behaviour, including a secular slow down accounted for (on average) by a cubic polynomial. Significant phase noise was observed in form of glitches, jumps and random walks in frequency, but their origin remains as yet unknown. From an investigation of the radio timing of a sample of pulsars (not including the Crab) it has been proposed that often this behaviour may be induced by the abrupt change between two different spin-down rates resulting from variations in the pulsars magnetosphere \\citep{lyne+10}.  Recent investigations focused on the delay between the arrival time of the main pulse between the radio and other wavelengths, including the optical (e.g. \\citealt{oosterbroek+06,oosterbroek+08}), and showed that the pulses in different energy bands are not aligned. This has profound implications for the geometry of the pulsar emission regions. All these studies have shown the relevance of independent timing studies in the optical (and other energy bands) to understand the pulsar emission mechanism by comparison with the radio. Our observations of the Crab pulsar with Iqueye are aimed at obtaining the most accurate optical timing needed to address these issues.  The plan of the paper is the following. Section~\\ref{sec2} presents the Iqueye observations of the Crab pulsar and the adopted data reduction procedure. In Section~\\ref{sec3} we show the phase fitting of the 2009 January and December data while, in Section~\\ref{sec4}, a comparison with the results from the radio ephemerides is presented. Section~\\ref{sec5} is devoted to the phase coherent phase fit of the whole year 2009 data (some mathematical details are reported in Appendix~\\ref{app}) and Section~\\ref{sec6} to some conclusive remarks. ",
        "conclusions": "\\label{sec6}  The measurement of the optical$-$radio phase shift, obtained with Iqueye in 2009 ($\\sim 240 \\, \\mu$s in January; $\\sim 160 \\, \\mu$s in December), is consistent with that determined in our previous measurement performed with Aqueye in Asiago in 2008 October ($\\sim 230 \\, \\mu$s; \\citealt{germana+12}). The optical pulse always leads the radio one. This is also in agreement with a number of other measurements performed after 1996 (see e.g. \\citealt{oosterbroek+08}). Only a couple of previous measurements yielded different results (consistent with no or negative delay; \\citealt{golden+00,romani+01}), possibly because some systematics were not adequately accounted for (e.g. the difference between the phase of the peak of the main pulse and the phase of the centroid of the main pulse, which is due to the asymmetric pulse shape). Therefore, although it is not possible to exclude a secular change of the phase delay, the available data show that the pulse in the optical leads that in the radio by 150-250$\\mu$s (see also \\citealt{strader+13}). This can be caused by a different position of the two emitting regions (50$-$80 km), although this is difficult to understand as electrons close to the pulsar are highly relativistic and any travel time delay between electrons and photons is essentially negligible. Most probably, the optical and radio beams are misaligned (1.5$^0$$-$3$^0$) because at the position where electrons emit optical photons the magnetic field has a slightly different orientation with respect to that where radio emission takes place.   The counting statistics available in the optical band allows us to compute the timing solution with extreme accuracy from a small number of consecutive observing nights, with an error of only $\\approx 3\\times10^{-5}$ in fractional phase or $\\approx 1 \\, \\mu$s in time (see Table~\\ref{tab3}). The optical timing solution matches the radio phase quite closely, yet in addition to a presumably constant phase shift between optical and radio, we find small $\\approx \\ 10 \\, \\mu$s d$^{-1}$ phase drifts between the two solutions, which are consistent with timing noise deduced from radio and optical data and with residuals effects of changing dispersion measures on the radio data.  The distributions of phase residuals are well approximated by Gaussians. The best observations have a dispersion $\\sigma \\approx 15 \\, \\mu$s (with 2 s sampling), in agreement with theoretical expectations for pure photon noise-induced phase variability. However, a few January and December observations have a slightly larger spread ($\\approx 20 \\, \\mu$s), which is most likely caused by worse quality of the sky. We note that the dispersion of residuals in our data is much smaller than the dispersion of the residuals for Giant Radio Pulses (84 $\\mu$s) recently reported by \\citet{rameshbhat+13}.     Giving the exquisite precision of the Iqueye data, we were able to phase connect the January and December observing runs, obtaining a global third-order timing solution for year 2009 and an independent measurement of the pulsar braking index. A comparison with the 2009 timing solution in the radio band shows that the pulsar makes small amplitude random excursions/oscillations around the ``average'' cubic spin-down (with residuals up to a few ms in phase and a jitter up to $\\sim 10^{-8}$ Hz in frequency), that are statistically significant. Imposing that the number of optical turns be equal to that in the radio band, the optical and radio phases in January and December are very close, but at intermediate times phase predictions differ. As a consequence, the inferred values of the frequency, its derivatives and braking indices also differ. Therefore, while the third-order secular spin-down and braking index model are very good at reproducing the overall evolution of the phases (e.g. \\citealt{groth75c}), an accurate description of the pulsar phase excursions/oscillations requires a frequent monitoring on a monthly time-scale.  The optical observations reported here confirm the third-order secular spin-down and braking index model to a high degree of accuracy. During the period from 2009 January to December, the pulsar has followed a braking index law so closely that, from the two Iqueye sets of observations alone, we were able to count the almost nine hundred million turns that the pulsar has made with an uncertainty of only two turns.  Future and more regular observations taken with both Iqueye and its twin Aqueye, mounted at the Copernico telescope in Asiago, will offer an invaluable means to monitor the optical$-$radio phase delay, to understand the phase excursions around the braking index model and to pinpoint the deep physical origin of the pulsar spin-down mechanism."
    },
    "1402/1402.0268_arXiv.txt": {
        "abstract": "The recent detection of Fast Radio Bursts (FRBs) has generated strong interest in identifying the origin of these bright, non-repeating, highly dispersed pulses. The principal limitation in understanding the origin of these bursts is the lack of reliable distance estimates; their high dispersion measures imply that they may be at cosmological distances ($0.1 < z < 1.0$). Here we discuss new distance constraints to the FRB010621 (a.k.a J1852$-$08) first reported by Keane. We use velocity resolved $H\\alpha$ and $H\\beta$ observations of diffuse ionised gas toward the burst to calculate an extinction-corrected emission measure along the line of sight. We combine this emission measure with models of Galactic rotation and of electron distribution to derive a 90\\% probability of the pulse residing in the Galaxy.  However, we cannot differentiate between the two Galactic interpretations of Keane: a neutron star with unusual pulse amplitude distribution or Galactic black hole annihilation. ",
        "introduction": "In the past few years, there have been a growing number of  detections of unusual radio pulses from the Parkes radio telescope, known as Fast Radio Bursts (FRBs). The six known FRBs \\citep{lorimer2007bmr, Keane12,Thornton13}, are characterised by extremely high flux densities ($\\sim$ Jy) over very short time scales (milli-seconds), high dispersion measures (DMs), and in one case, a scatter-broadened pulse profile. While there was for some time doubt as to whether these pulses are astronomical or terrestrial \\citep{BurkeSpolaor11lg, Kocz12},  more recent discoveries support an astronomical interpretation \\citep{Keane12, Thornton13}, in spite of the lack of multi wavelength counterpart.   It has not been definitively shown if the pulses are Galactic or extragalactic. Most of these discoveries favour an extragalactic origin ($0.1 \\lesssim z \\lesssim 1.0$), primarily because the DM substantially exceeds both the known dispersion measures from pulsars near to the line of sight, and / or a DM obtained from a commonly used Galactic model (e.g. NE2001 \\citet{Cordes02}). If they are extragalactic, their extremely high luminosities can only be explained by exotic mechanisms such as annihilating black holes \\citep{Rees77}, binary neutron star mergers \\citep{Hansen01, Totani13} `combing' magnetic fields in expanding supernova shells \\citep{Colgate71}, magnetar bursts \\citep{Popov13}, collapse of supramassive neutron stars \\citep{Falcke13} or gamma-ray bursts \\citep{Zhang13, Bannister12}. There is considerable interest in identifying such classes of objects; they are also prime candidates as the electromagnetic counterparts to high frequency gravity wave sources that are expected to be discovered with upcoming facilities \\citep{LIGO2012, Nissanke13}.  However, if fast radio bursts are Galactic, their characteristics can be attributed to somewhat less exotic mechanisms such as pulsar giant pulses with unusual pulse amplitude distributions \\citep{Keane12}, or nearby flaring stars, where the large dispersion measure is explained by a blanket of ionised material around the star \\citep{Loeb13}.  Because the progenitors of all of these highly-dispersed radio bursts are unknown, and because there are a small number of detections, it is important to examine each of the detections in detail. In this paper, we aim determine whether the burst named FRB010621 resides in the Galaxy or not.  FRB010621 (also known as J1852$-$08), was first reported by \\citet{Keane11} and discussed in more detail by \\citet{Keane12}. The burst position happens to lie toward an unusually low extinction region of the Galaxy just off the plane ($b=-4^\\circ$).  Optical emission line spectroscopy of diffuse ionised gas has been studied extensively in this region \\citep{Madsen05}. This fortunate coincidence allows us to use this spectroscopy to estimate the amount of Galactic ionised material in the direction of  FRB010621, and therefore determine the extent to which the DM of  FRB010621 (746 \\pccm) can be attributed to ionised gas in our own Galaxy. The use of such spectroscopy has been used to successfully model distances to Galactic pulsars, albeit with less optimal dust corrections \\citep{Schnitzeler12}. Our distance constraint will aid in the understanding the progenitor of  FRB010621 and in understanding statistical properties of other highly-dispersed radio bursts.  This paper is organised as follows: in \\S\\ref{sec:scutum} we describe existing observations of the low-extinction region of the ISM in which FRB010621 falls.  We discuss the emission measure toward the burst in \\S\\ref{sec:em} and outline a model that relates the EM to the DM in \\S\\ref{sec:dmcalc}.  In \\S\\ref{sec:distance} we describe our constraints on the distance to the burst. We discuss the implications of our results on the origin of FRB010621 in \\S\\ref{sec:discussion} and summarise our results in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have taken advantage of the fortuitous location of the fast radio burst FRB010621 toward a low-extinction region of the Galaxy to estimate the extinction corrected emission measure toward the burst. We have combined this with the observed DM to the pulse and evaluated the pulse distance and DM to the edge of the Galaxy, and their uncertainties, with a Monte Carlo simulation. Under the assumption that the interstellar dust toward the pulse acts as a foreground screen, we find  that FRB010621 is located within the Galaxy with 90\\% confidence and the distance is $14\\pm6$~kpc. While our results strongly favour a Galactic interpretation, we cannot differentiate between the two Galactic scenarios proposed by Keane: namely a pulsar with unusual amplitude distribution, or annihilating Galactic black hole.  Our results also show the limitations of using the Galactic electron models, such as NE2001 to estimate distances to dispersed pulses, especially near the Galactic plane.  We note that while FRB010621 is likely to reside in the Galaxy, other FRBs may well be extragalactic. The other FRBs are at much higher Galactic latitude where the electron column density is not large enough to account for the observed DMs by much larger factors.  It therefore could be that FRB010621 belongs to a different class of FRBs. There are a number of archival and new observational surveys that are searching for FRBs \\citep{Macquart10, Hassall13, Trott13}.  We anticipate that the discovery of more of these kinds of bursts will significantly increase our limited understanding of this strange phenomenon."
    },
    "1402/1402.7354_arXiv.txt": {
        "abstract": "The stellar initial mass function (IMF) is a key parameter to study galaxy evolution. Here we measure the IMF mass normalization for a sample of 68 field galaxies in the redshift range 0.7 to 0.9 within the Extended Groth Strip. To do this we derive total (stellar + dark matter) mass-to-light [$(M/L)_{\\rm dyn}$] using axisymmetric dynamical models. Within the region where we have kinematics (about one half-light radius), the models assume: (i) that mass-follows-light, implying negligible differences between the stellar and total density profiles; (ii) constant velocity anisotropy ($\\beta_{\\rm z}\\equiv1-\\sigma_z^2/\\sigma_R^2=0.2$); (iii) that galaxies are seen at the average inclination for random orientations (i.e. $i=60^\\circ$, where $i=90^\\circ$ represents edge-on). The dynamical models are based on anisotropic Jeans equations,  constrained by HST/ACS imaging and the central velocity dispersion of the galaxies, extracted from good-quality spectra taken by the DEEP2 survey. The population $(M/L)_{\\rm pop}$ are derived from full-spectrum fitting of the same spectra with a grid of simple stellar population models. Recent dynamical modelling results by the \\atl\\ project and numerical simulations of galaxy evolution indicate that the dark matter fraction within the central regions of our galaxies should be small. This suggest that our derived total $(M/L)_{\\rm dyn}$ should closely approximate the stellar $M/L$. Our comparison of $(M/L)_{\\rm dyn}$ and $(M/L)_{\\rm pop}$ then imply that for galaxies with stellar mass $M_\\ast\\ga10^{11}$ \\msun, the {\\em average} normalization of the IMF is consistent with a Salpeter slope, with a substantial scatter. This is similar to what is found within a similar mass range for nearby galaxies. ",
        "introduction": "Knowledge of the stellar mass of galaxies has long been pivotal to the study and testing of galaxy formation and evolution theories, but an economical method for it's estimation has remained elusive to astronomers for decades \\citep[see][for a review]{Courteau2014}. The main reason for this is the fact that the galaxy photometry doesn't sufficiently constrain the stellar mass, which is dependent on both the stellar populations of the galaxy and the stellar Initial Mass Function (IMF) of those populations. The IMF represents the distribution of stellar masses at a single star formation event. The form of the IMF is critical for estimating the mass of the galaxies and in understanding the stellar feedback and the chemical enrichment processes taking place within them.  It is currently accepted, from direct star counts, that the IMF shape in our Milky Way can be approximated by a power-law $d\\,N/d\\,m \\propto m^x$ with the \\citet{salpeter1955} slope $x\\approx-2.35$, above a stellar mass $m\\ga0.5$ \\msun\\ and a more shallow slope below $m \\la 0.5\\msun$ \\citep{kroupa2001,chabrier2003}. Though this IMF seems universal within the Milky Way its applicability to external galaxies has been uncertain for decades (see \\citealt{Bastian2010} and \\citealt{kroupa2013} for recent reviews).  Early attempts, comparing the dynamical mass-to-light ratio M/L to the value inferred from stellar population synthesis, concluded that at least some spiral galaxies \\citep{bell2001,kassin2006,bershady2011,Dutton2011,brewer2012} and ETGs \\citep{cappellari2006,Ferreras2008,Thomas2011} require an overall Kroupa/Chabrier mass normalization of the IMF like the Milky Way. By studying the variation of the average M/L of galaxies in clusters at different redshifts, \\cite{renzini2006} concluded that a Salpeter IMF slope was required for the mass range of $1<M<1.4$ \\msun.  More recently \\citet{vandokkum2010,vanDokkum2011} studied near-infrared spectral features, where the contribution from dwarf stars is easier to measure. They concluded that massive elliptical galaxies have a dwarfs-dominated IMF, unlike the Milky Way. \\citet{auger2010b} studied the IMF by comparing stellar masses from strong gravitational lensing to stellar masses from population synthesis. Their models assume massive ETGs are spherical, satisfy a power law dependence of IMF normalization with galaxy mass, have dark halos following the mass-concentration relation predicted by simulation and the total to stellar mass derived by halo abundance matching. They concluded massive ETGs as a class have a Salpeter-like mass normalization of the IMF.  \\citet{cappellari2012nature} used more general axisymmetric dynamical models allowing for the IMF, the stellar profiles, the halo slope and the stellar fraction to freely vary in different galaxies to reproduce the data. The models were constrained by integral-field stellar kinematics for a sample of 260 ETGs spanning a large range of masses and $\\sigma$. They found a systematic trend in the IMF normalization. The average IMF varies from Kroupa/Chabrier to Salpeter and heavier with increasing $\\sigma$, which was shown to trace the bulge fraction \\citep{atlas3d20}.  A number of recent works appear so far consistent with a systematic IMF variation. These used either spectral synthesis to study absorption features \\citep[e.g.][]{Spiniello2012,Conroy2012,Smith2012,Ferreras2013,LaBarbera2013}, dynamical scaling relations \\citep{Dutton2013} or approximate spherical dynamical models \\citep{tortora2013}. However quantitative agreement between the different approaches has not yet been achieved and full consensus has not been reached \\citep[e.g.][]{Maraston2013}.  The study of IMF in galaxies at high redshift is so far dominated by dynamical mass derivations using the virial estimator \\citep[e.g.][]{renzini2006,vandeSande2013}. In this study, we constrain the IMF normalization at redshift $z\\sim1$ using axisymmetric dynamical models and compare the results with findings in the local universe. This is one of the few studies to attempt dynamical modelling at higher redshift \\citep{vanderMarel2007,cappellari2009} but is the first to place constraints on the IMF normalization of galaxies at high redshift.  In section 2, we give a description of the data used for this study and a description of the galaxy set. We describe our methodology in section 3, while we present our results in section 4. In section 5, we discuss the results. We used the following cosmological constants: $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, $H_{\\rm 0}=70$ \\kms\\ Mpc$^{-1}$. ",
        "conclusions": "We used the JAM modeling method to derive dynamical $(M/L)_{\\rm JAM}$ for 68 massive galaxies ($M_\\ast=10^{11}-10^{12}$ \\msun) at redshifts $z=0.7-0.9$. Our results show that using the virial equation at high redshifts can lead to severe under-estimation in the dynamical masses of galaxies, which is likely due to the under-estimation of the $R_{\\rm e}$. The under-estimation must depend on the depth of the photometry as well as on the adopted extrapolation of the profiles, making the virial estimates assumption dependent \\citep[see also][]{atlas3d15}. The $(M/L)_{\\rm JAM}$ doesn't suffer for this problem as it only requires the knowledge of the light and mass distribution within a region comparable to where the stellar kinematics are available.  Our result suggests that stellar population estimates of stellar masses of massive galaxies at high redshift should assume a Salpeter IMF normalization for more accurate results, instead of the Kroupa/Chabrier IMFs which are often adopted. This IMF normalization is the same inferred for nearby galaxies of similar masses."
    },
    "1402/1402.0902.txt": {
        "abstract": "{Planet-forming disks of gas and dust around young stars contain polycyclic aromatic hydrocarbons (PAHs). % PAHs can exist in different charge states depending on the physical properties of the star and surrounding disk. } %Aims {We aim to characterize how the charge state of PAHs can be used as a probe of flows of gas through protoplanetary gaps. In this context, our goal is to understand the PAH spectra of four transitional disks. In addition, we want to explain the observed correlation between PAH ionization (traced by the $I_{6.2}/I_{11.3}$ feature ratio) and the disk mass (traced by the 1.3 mm luminosity). } %Methods { We implement a model to calculate the charge state of PAHs in the Monte Carlo radiative transfer code MCMax. The emission spectra and ionization balance are calculated in the parameter space set by the properties of the star and the disk. } %Results { A benchmark modeling grid is presented that shows how PAH ionization and luminosity behave as a function of star and disk properties. The PAH ionization is most sensitive to ultraviolet (UV) radiation and the electron density. In optically thick disks, where the UV field is low and the electron density is high, PAHs are predominantly neutral. Ionized PAHs trace low-density optically thin disk regions where the UV field is high and the electron density is low. Such regions are characteristic of gas flows through the gaps of transitional disks. We demonstrate that fitting the PAH spectra of four transitional disks requires a contribution of ionized PAHs in `gas flows' through the gap. %There is no gradient in the PAH charge state as a function of disk radius in an optically thick disk. } %Conclusions { The PAH spectra of transitional disks can be understood as superpositions of neutral and ionized PAHs. For  HD\\,97048, neutral PAHs in the optically thick disk dominate the spectrum. In the cases of HD\\,169142, HD\\,135344\\,B and Oph IRS 48, small amounts of ionized PAHs located in the `gas flows' through the gap are strong contributors to the total PAH luminosity. The observed trend in the sample of Herbig stars between the disk mass and PAH ionization may imply that lower-mass disks have larger gaps. Ionized PAHs in \u00d4gas flows\u00d5 through these gaps contribute strongly to their spectra. } ",
        "introduction": "%Many astrophysical objects show a rich spectrum of infrared (IR) emission features associated with Polycyclic Aromatic Hydrocarbons (PAHs, \\citealt{1984LegerPuget,1985Allamandola}). PAH features are ubiquitously found in space, examples include HII regions, reflection nebulae, young stellar objects, planetary nebulae, post-asymptotic giant branch objects, nuclei of galaxies, and ultra-luminous infrared galaxies (\\citealt{2008Tielens} and references therein). PAHs account for 10\\%-30\\% of cosmic carbon \\citep{1995SnowWitt, 1989PugetLeger, 1989Allamandola}. They have been detected in cometary material during the Deep Impact mission \\citep{2006Lisse} and were brought back to Earth by the Stardust mission \\citep{2006Sandford}. It is therefore believed that PAHs are also important contributors to the composition of planets in habitable zones \\citep{1993Kasting}.  The infrared (IR) spectra originating from protoplanetary disks are rich and diverse. Polycyclic aromatic hydrocarbons (PAHs) are observed in disks around Herbig Ae/Be stars \\citep{2001Meeus, 2004AckeAncker, 2008Keller, 2010Acke} and T Tauri stars \\citep{2007bGeers, 2008Bouwman}. The strength of the PAH features decreases with stellar effective temperature. No T Tauri stars of the spectral type later than G8 shows the IR emission features in their spectra  (e.g. \\citealt{2008Tielens,2011Kamp} and references therein). In addition, the strength of the IR emission features is generally much weaker relative to the total IR emission in the Herbig star spectra as compared to the diffuse ISM  \\citep{2005Sloan, 2008Keller, 2010Acke, 2008Boersma}. Direct imaging observations taken in the PAH bands have shown that PAHs can be used as a tracer of the outer disk surface \\citep{2006Doucet, 2006Lagage}.  Even so, the PAH emission is not dominated by PAHs on the surface of the outer disk in some transitional disks \\citep{2007aGeers, 2013Maaskant}.  After being electronically excited by ultraviolet (UV) photons, the PAHs cool by emission in the CH- and CC- stretching and bending modes. Laboratory studies and quantum chemical calculations show that many of the IR features shift in peak position, vary in width, and/or show substructure depending on ionization stage and detailed molecular structure \\citep{1999HudginsAllamandola,1999Allamandola, 2009Bauschlicher, 2012Ricca}. An important parameter that influences the relative feature strength of the CH and CC modes is the effect of ionization with CC modes being carried predominantly by ions and CH modes by neutrals. Hence, the 6.2-$\\upmu$m (CC--mode) to 11.2-$\\upmu$m (CH--mode) ratio provides a good measure of the ionization balance of the emitting PAHs \\citep{1999Allamandola}. This has been demonstrated in a variety of astrophysical environments \\citep{2001Hony, 2002Peeters, 2005Rapacioli, 2008Galliano}.   \\begin{figure*}[htbp] \\includegraphics[width=0.52\\textwidth]{plot_normalized.pdf} \\includegraphics[width=0.22\\textwidth]{FWHM_Nband_all_FWHM.pdf} \\includegraphics[width=0.25\\textwidth]{fourobjects_v2.jpg} \\caption{ \\label{four_objects_PAHs} Left: Spitzer/IRS spectra of the four transitional disks HD\\,97048, HD\\,169142, HD\\,135344\\,B and Oph IRS 48. The spectra are scaled relative to the 6.2-$\\upmu$m PAH peak flux and show a wide range in the $I_{6.2}/I_{11.3}$ ratio; most neutral in HD\\,97048 and most ionized in Oph IRS 48.  Middle: VLT/VISIR observations show that FWHM as a function of wavelength. The size of the 11.3-$\\upmu$m feature is respectively larger, similar and smaller than the N-band continuum for HD\\,97048, HD\\,169142, HD\\,135344\\,B, and Oph IRS 48 (data from \\citet{2013Maaskant}). The diamond symbols show the sizes of the point spread functions of the observations. The black dashed-dotted line is the diffraction limit of the telescope. Right: Sketches showing a qualitative interpretation of the spatial distribution of the dominant contribution of PAH emission, as shown by the colored ovals. Each object is given a similar color on the left, middle, and right plots. %A low $I_{6.2}/I_{11.3}$ ratio is indicative of neutral PAHs from the disk (e.g. HD\\,97048), while a high $I_{6.2}/I_{11.3}$ ratio traces ionized PAH emission originating from the disk gap (e.g. Oph IRS 48).} } \\end{figure*}    The intense radiation field of the central star may photo-dissociate the PAH molecules. As a result, the average PAH size increases, and the total abundance decreases \\citep{2007Visser, 2012Siebenmorgen}. As a possible path for PAH-survival, \\citet{2010Siebenmorgen} suggest that turbulent motions in the disk can replenish or remove PAHs from the reach of hard photons. Several modeling studies have been carried out to understand global trends in the composition and evolution of PAHs in relation to the disk \\citep{2004Habart, 2007bDullemond, 2009Berne}. However, none of them are focused on understanding PAH ionization.  In this paper, we investigate the effect of disk environments on PAH charge, spectra, and luminosity. The key question is: can we use the ionization balance of PAHs as a tracer of processes in protoplanetary disks? In Section \\ref{sec:obstrends}, we first identify new observational trends seen in the $I_{6.2}/I_{11.3}$ PAH feature strength ratios. Section \\ref{sec:PAHmodel} describes the PAH model in the Monte Carlo radiative transfer tool MCMax. In Section \\ref{sec:results}, we show how the ionization of PAHs behaves as a function of star and disk properties. We use these results to demonstrate how we can understand the PAH spectra of four transitional disks in Section \\ref{sec:PAHfits}. In the discussion (Section \\ref{sec:discussion}), we discuss what we can learn about disk evolution from PAH ionization. The conclusions are given in Section \\ref{sec:conclusions}.  %----------------------------------------------------------------------------------------- %\t\t\t\t\tOBSERVATIONS %----------------------------------------------------------------------------------------- ",
        "conclusions": "\\label{sec:conclusions} We have modeled the charge state of PAHs in Herbig disks. We find that the neutral fraction of the disk is high for typical star and disk parameters. The key parameter to ionize the PAHs in the disk is to lower the gas density. The PAH luminosity is most influenced by competition from dust grains, either by an increased PAH mass fraction, or the removal of small dust grains. Our results can be summarised as follows:  \\begin{itemize} \\item Emission from ionized PAHs trace low-density optically thin regions and has a high $I_{6.2}/I_{11.3}$ ratio. \\item The PAHs in an optically thick disk are predominantly neutral and produce emission with a low $I_{6.2}/I_{11.3}$ ratio. \\item The PAH spectra of transitional disks can be understood as superpositions of neutral PAHs from an optically thick disk and ionized PAHs from that of low density, optically thin gas flows through protoplanetary disk gaps. \\item There is no gradient in the neutral fraction of PAHs as a function of radius in an optically thick disk. \\item The PAH luminosity is not synonymous to the disk structure. As shown for the transitional disks in our sample, the total PAH mass in the low density, optically thin disk gap can be much smaller compared to the total PAH mass in the outer disk. Nevertheless, they give a dominating contribution to the total PAH luminosity. \\item We find a trend in the $I_{6.2}/I_{11.3}$ ratio compared to the 1.3 millimeter luminosity. Since the 1.3 millimeter luminosity can be used as a proxy for the disk mass, the degree of ionization tends to be higher for lower disk masses. We suggest that the contribution to the PAH luminosity from ionized PAHs in low-density optically thin `gas flows' becomes larger for lower mass disks. \\item For a sample of 18 Herbig stars, the $I_{6.2}/I_{11.3}$ ratios indicates that PAH emission from flaring/transitional disks (group I) is higher and originates in a wider range of physical environments  \\mbox{($\\langle I_{6.2}/I_{11.3} \\rangle = 2.61 \\pm 1.63$)}, while the origin of PAH emission in self-shadowed flat disks (group II) seems to be more homogeneous \\mbox{($\\langle I_{6.2}/I_{11.3} \\rangle = 2.08 \\pm 0.56$)}.  \\end{itemize}  With superior mid-IR sensitivity and spatial resolution, observations by the Mid-InfraRed Instrument (MIRI) on the James Webb Space Telescope (JWST) will be able to probe more and fainter targets. This work shows that the ionization balance of PAHs  provides an extra diagnostic tool to characterize the gaps of transitional disks and, thus, to study the process of planet formation."
    },
    "1402/1402.2938_arXiv.txt": {
        "abstract": "We present updated magnetic field maps of the chemically peculiar B9p star $\\alpha^2$ CVn created using a series of time resolved observations obtained using the high resolution spectropolarimeters ESPaDOnS and Narval. We compare these new magnetic field maps with the original magnetic Doppler imaging maps based on spectra recorded with the MuSiCoS spectropolarimeter and taken a decade earlier. These new maps are inferred from line profiles in all four Stokes parameters using the magnetic Doppler imaging code INVERS10. With the addition of new lines exhibiting Stokes $IQUV$ signatures we have a unique insight into how the derived magnetic surface structure may be affected by the atomic lines chosen for inversion.  We report new magnetic maps of  $\\alpha^2$ CVn created using strong iron lines (directly comparable to the published MuSiCoS maps), weak iron lines and chromium lines, all of which yield a magnetic field structure roughly consistent with that obtained previously.  We then derive an updated magnetic structure map for $\\alpha^2$ CVn based on the complete sample of Fe and Cr lines, which we believe to produce a more representative model of the magnetic topology of $\\alpha^2$ CVn. In agreement with the previous mapping, this new updated magnetic map shows a dipolar-like field which has complex sub-structure which cannot be explained by a simple low order multipolar geometry.  Our new maps show that regardless of the atomic line or species choice, the reconstructed magnetic field is consistent with that published previously, suggesting that the reconstructed field is a realistic representation of the magnetic field of $\\alpha^2$ CVn.  $\\alpha^2$ CVn is the first Ap star for which multiple, high resolution magnetic maps have been derived, providing important observational evidence for the stability of both the large and small-scale magnetic field. ",
        "introduction": "The bright Ap star $\\alpha^2$~CVn has been the subject of many observations over the past century, with the first period determination as early as Markov (1930), followed by Farnsworth (1932). There have been many studies of the magnetic and spectral variability of  $\\alpha^2$~CVn (Babcock \\& Burd (1952); Pyper (1969); Borra \\& Landstreet (1977)). It was not until Kochukhov et al. (2002) employed a new magnetic Doppler imaging technique (MDI) (described by Piskunov \\& Kochukhov (2002) and Kochukhov \\& Piskunov (2002)) that the first high resolution maps of the surface vector magnetic field using Stokes $IV$ observations were made for $\\alpha^2$~CVn. These maps were later refined by using linear polarisation profiles (Stokes $Q$ and $U$) in combination with Stokes $IV$ (Kochukhov \\& Wade 2010) acquired with the MuSiCoS spectropolarimeter. These maps (along with those of 53 Cam, Kochukhov et al. 2004) were distinguished from earlier models in that they were computed directly from the observed polarised line profiles, making no {\\em a priori} assumptions regarding the large-scale or small-scale topology of the field. The MDI surface magnetic field maps of both stars revealed that their magnetic topologies depart significantly from low-order multipoles.  These original maps were limited by the quality of the observational data. With MuSiCoS being a relatively inefficient medium resolution spectropolarimeter, only a very small number of lines exhibited Stokes $QU$ profiles of sufficient quality for modelling. With the new observations of Ap stars in all four Stokes $IQUV$ parameters obtained using the new ESPaDOnS and Narval instruments as described by Silvester et al. (2012), it is now possible to not only study more spectral lines, these new data also allow the study of $\\alpha^2$~CVn at a resolution not previously possible and allow a more detailed probe of subtle spectral features which have been unresolved or buried in the noise in the MuSiCoS observations.  The mapping performed by Kochukhov \\& Wade (2010) (herein referred to as K\\&W) used the inversion code INVERS10, which is also used in this study. INVERS10 employs a single mean metallicity model atmosphere when performing the inversions. It was suggested by Stift et al. (2012) that using a single mean model atmosphere, which did not account for horizontal (local) atmospheric variations would lead to the derivation of incorrect abundance distributions and incorrect magnetic field geometries. To address these concerns, Kochukhov et al. (2012) compared magnetic field maps of $\\alpha^2$~CVn reconstructed with INVERS10 with those reconstructed using a version of the INVERS code which incorporated horizontal variation of the model atmosphere. Kochukhov et al. (2012) found no significant differences between the mapping results from the two codes, confirming the suitability of INVERS10 for mapping both the magnetic field and chemical abundance features in Ap stars such as $\\alpha^2$~CVn. This work aims to further explore the results of Kochukhov at al. (2012), by investigating whether the same {\\it unique} magnetic field topology can be obtained from various sets of suitable Stokes $IQUV$ lines, taken with the higher spectral resolution data from ESPaDOnS and Narval.  \\begin{table} \\caption{Fundamental parameters used/derived for  the $\\alpha^2$ CVn mapping.  {\\bf References}: (1) Kochukhov et al. (2002), (2) Farnsworth (1932), (3) Kochukhov and Wade (2010)} \\begin{tabular}{ccc} \\hline \\hline Parameter & Value & Reference \\\\ \\hline \\teff &  $11600 \\pm 500$ K  & (1) \\\\ \\lgg & $3.9 \\pm 0.1$ & (1) \\\\ $P_{rot}$ & 5.46939 days & (2) \\\\ \\vs & $18.0 \\pm  0.5$ \\kms& \\\\ $i$ & $120^{\\circ} \\pm 5$ & (3) \\\\ $\\Theta$ & $115^{\\circ} \\pm 5$ & (3) \\\\ \\hline \\label{parameter-table} \\end{tabular} \\end{table}  The remainder of the paper is organised as follows: section 2 briefly describes the observations, section 3 discusses the procedure for selecting lines suitable for inversion. In section 4 we discuss the various magnetic maps of $\\alpha^2$~CVn, with the results and implications of these maps. Finally we summarise our findings and the implications in the conclusion. ",
        "conclusions": "With the results presented in section 4.1 we are able to show that the magnetic maps from K\\&W and magnetic maps created using the new observations of Silvester et al. (2012) are indeed consistent when considering the iron lines Fe~{\\sc ii}\\ $\\lambda$ 4923 and 5018, with only small differences likely attributable to the difference in data quality (spectral resolution and signal-to-noise) between the two data sets. This also illustrates that the small scale structure is stable over the timescale between the two epochs of observation.  In section 4.2 we investigated the potential impact of line selection on the resulting magnetic map by using only weak iron lines in the inversion.  This produced a magnetic map that was broadly consistent with the strong line map and only differed in the meridional field at low latitudes and the azimuthal field at high latitudes. The largest difference in the azimuthal field is on the order of 1 kG and located around a latitude of $25^{\\circ}$. The difference in the meridional field is on the order of 1.5 kG and located around a latitude of $-60^{\\circ}$. Both differences are located around a longitude of $180^{\\circ}$. This difference is believed to be a result of the difference in horizontal structure sensitivities between the weak and strong line sets, with stronger saturated lines altered less by abundance spots. It is important to note that profiles computed from weak lines did not agree well with profiles computed from strong iron lines. In addition the meridional field is the most difficult to constrain by inversion. This suggests the strong iron map represents a smoothed version of the magnetic field map.  It should be noted that a potential source of the discrepancy seen between the weak and the strong iron magnetic maps is vertical stratification.  Theoretical modelling of abundance stratification by Alecian and Stift (2010) showed that in an Ap star, vertical stratification leads to a change in chemical abundance as a function of optical depth in the atmosphere. In such a framework, strong lines which sample a larger range of optical depths when compared to weak lines, would be probing different parts of this ``abundance vs optical depth'' variation.  This effect is more significant in cooler Ap stars,  however for  hotter stars like $\\alpha^2$ CVn,  Stift and Alecian (2012) showed it could potentially lead to variations in abundances on the order of 1 dex over the range of optical depths.  By combining strong and weak iron lines the resulting ``combined iron line'' map was slightly more consistent with the strong iron line map, than when weak iron lines alone were used in the mapping. The contribution of the weak lines into the combined iron line map still results in a difference in the meridional field at low latitudes. It was possible to reproduce both the weak and strong iron line profiles with this single map.  In section 4.4 we produced a magnetic map based on chromium lines. The results were similar to those found by mapping weak iron lines, with an overall agreement with the strong iron line map but still with differences in the meridional fit at low latitudes.  Interestingly both the weak iron line map and the chromium lines independently give very similar magnetic field maps. This illustrates that very different line sets can give consistent mapping results.  Finally in section 4.6 we combined all the above line sets into a final map of the magnetic field of $\\alpha^2$ CVn. We consider this map to be the best existing representation of the magnetic field of $\\alpha^2$ CVn. The overall structure of the final map is in agreement with K\\&W: we find a dipole-like structure with complex sub-structure. As was found by K\\&W, we find that the magnetic field is strongest at the positive pole (seen clearly at phase 0.6), with an asymmetry compared to the negative pole.  There is clear agreement between the maps of K\\&W and our new maps (produced from a completely new set of Stokes $IQUV$ observations). When this is combined with the fact that we are able to reproduce the same general magnetic field topology from a variety of lines sets, of varying intensities and from two different atomic species, it provides compelling support to the findings of Kochukhov et al. (2012), and suggests the magnetic field structure we reconstruct for  $\\alpha^2$ CVn is accurate and not the result of limitations in the inversion technique as has been suggested by Stift et al. (2012). In addition we have also illustrated that the observed profiles cannot be fit with a simple dipolar or dipole + quadrupole geometry and can only be fit with direct magnetic mapping using only local regularization to constrain the map. One other important result from this work is that we have the first confirmation via magnetic Doppler imaging that the global magnetic field of $\\alpha^2$ CVn is stable over the period of a decade, which adds further evidence to the current theoretical understanding of the stability of the magnetic fields in Ap stars.  It should be noted that regardless of the line-set choice, the large-scale structure of all the maps is consistent with the original maps of K\\&W and strong iron line maps. Line selection may have an subtle effect on the resulting magnetic field map, with small differences arising as a result of horizontal structure sensitivity of the atomic line used in the inversion. With this in mind, it is clearly of value to map a variety of lines with different formation heights when possible, such as weak and strong iron lines  or strong iron lines and chromium lines, etc.  Even when it is not possible to have such a variety,  provided that the line profiles used are of sufficient quality and show clear linear polarisation amplitudes, the reconstructed field should be reliable.  It could be argued that data quality is a far bigger factor potentially affecting the reconstructed magnetic topology.  The next stage in the project is to produce a series of abundance maps for $\\alpha^2$ CVn, looking for any interplay between the magnetic field and chemical atmospheric structures.  We will then investigate the magnetic field geometry and chemical abundance structures of other Ap stars using data from Silvester et al. (2012) and the MiMeS project.  By increasing the sample of Ap stars studied using MDI, we can probe what influences other stellar parameters (such as mass, temperature, rotation etc) have on the resulting magnetic field geometry."
    },
    "1402/1402.2279_arXiv.txt": {
        "abstract": "{Understanding the history of formation of $z > 6$ quasars represents a major challenge to theoretical models. Physical insights on the connection between the central black hole and its host galaxy can be gained by means of the quasar infrared properties. } {Here we investigate the origin of the far-infrared continuum of SDSS J1148+5251, using it as a prototype for the more general class of high-luminosity high-redshift quasars. } {We run the radiative transfer code TRADING to follow the transfer of radiation from the central source and from stellar sources through the dusty environment of the host galaxy. We adopt simple models for the central source, including all the radiation that can travel beyond the dusty torus. The radiation from stellar sources is modeled using the code P{\\'E}GASE. The model is based on the output of the semi-analytical merger tree code, \\textsc{GAMETE/QSOdust}, which enables to predict the evolution of the host galaxy and of its nuclear black hole, following the star formation history and chemical evolution -- including dust -- in all the progenitor galaxies of SDSS J1148+5251. } {We find that the radiation emitted by the central source, which dominates the observed spectral energy distribution from UV/optical to near and mid infrared wavelengths, can also provide an important source of heating for the dust distributed in the host galaxy, powering at least 30\\% and up to 70\\% of the observed far infrared emission at rest-frame wavelengths $[20 - 1000]\\, \\mu$m. The remaining fraction is contributed by stellar sources and can only be achieved if the host galaxy is able to sustain a star formation rate of $\\rm \\approx 900 \\, M_{\\odot}/yr$ at $z=6.4$. This points to a co-evolution scenario where, during their hierarchical assembly, the first super-massive black holes and their host galaxies first grow at the same pace until the black hole reaches a mass of $\\rm \\sim 2 \\times 10^8 M_\\odot$ and starts growing faster than its host, reaching the bright quasar phase when the black hole and stellar mass fall within the scatter of the scaling relation observed in local galaxies. This same evolutionary scenario has been recently shown to explain the properties of a larger sample of $5 < z <6.4$ QSOs, and imply that current dynamical mass measurements may have missed an important fraction of the host galaxy stellar mass. } {We conclude that the far-infrared luminosity of high-z quasars is a sensitive tracer of the rapidly changing physical conditions in the host galaxy. Quasars appear far-infrared bright when the host galaxy can still sustain strong starbursts, with star formation rates $\\rm 100 \\, M_\\odot/yr < SFR < 1000 \\, M_\\odot/yr$ and progressively dim as large quasar-driven outflows deplete the host galaxy of its gas content, damping star formation and leaving the central source as the only source of dust heating.} ",
        "introduction": "\\label{sec:intro}   The first comprehensive far-infrared (FIR) investigations on optically selected low-$z$ AGNs have been made using {\\it Infrared Astronomical Satellite} ({\\it IRAS}) and {\\it Infrared Space Observatory} ({\\it ISO}) data, supplemented with SCUBA and IRAM data at longer wavelengths (see e.g. Haas et al. 2000, 2003). A variety of different Spectral Energy Distribution (SED) shapes was observed, suggesting that the mid-infrared (MIR) and FIR emission originate from dust components whose heating mechanisms may be different (Haas et al. 2000). In particular, while in all sources the MIR emission could be powered by the central Active Galactic Nuclei (AGN), the relative role of starbursts and AGN in powering the FIR emission was hard to established, although Haas et al. (2003) proposed that QSOs with the largest FIR luminosities ($\\rm L_{FIR} \\ge 10^{13} - 10^{14} L_\\odot$) favour a strongly dominating AGN contribution.  Progress in the field was enabled by observations with the {\\it Spitzer Space Telescope}, using MIR polycyclic aromatic hydrocarbon (PAH) emission features as tracers of star formation in QSO host galaxies (Schweitzer et al. 2006; Netzer et al. 2007;  Lutz et al. 2008). Indeed, the main MIR PAH emission features at 6.2, 7.7, 8.6, and 11.3 $\\rm \\mu m$  have been extensively used as extragalactic star formation indicators because of their presence in star forming environments with a wide range of physical conditions and their absence in the proximity of AGNs. While the calibration between the PAH luminosity and the star formation rate depends on the physical properties of the interstellar medium, PAH emission is sufficiently strong to be less easily outshone by the AGN emission, unlike $\\rm H\\alpha$ or the 24 $\\mu$m  continuum. Netzer et al. (2007) applied this technique to local QSOs; comparing the average SEDs of groups of QSOs that are differing in their FIR luminosity, they were able to derive a SED for the pure AGN, after subtraction of the host star formation. Lutz et al. (2008) have extended the study to a sample of mm-detected QSOs at $1.8 \\le z \\le 2.8$, showing that these sources follow the same correlation between the PAH luminosity and the rest-frame FIR emission inferred for local QSOs. For mm-bright QSOs, this provides strong evidence for intense star formation rates in their hosts, sometimes exceeding 1000 $\\rm M_\\odot/yr$. However, the mm-faint end of this QSO population may contain pure AGNs, with no strong host star formation, and the observed weak rest frame FIR emission appears consistent with pure AGN-heated dust (Lutz et al. 2008).  The launch  of the {\\it Herschel Space Observatory}  has enabled to directly observe the galaxy rest-frame FIR emission. Using large samples of type 1 and type 2 AGN in the {\\it Herschel} Multi-tiered Extragalactic Survey (HerMES) fields observed with SPIRE at  $\\rm 250, 350, \\, and \\,500 \\, \\mu$m, Hatziminaoglou et al. (2010) showed that quasars out to $z \\sim 4$ have the same FIR colours as star-forming galaxies and that a starburst was always needed to reproduce the FIR emission (see also Elbaz et al. 2010). By matching a sample of 24\\,$\\mu$m - selected QSOs in the {\\it Spitzer} Wide-area InfraRed Extragalactic Survey to HerMES, Dai et al. (2012) collected 32 QSOs with $0.5 \\le z < 3.6$ with SPIRE data. These FIR-detected QSOs show cold dust emission with $\\rm L_{FIR}$ in the range $\\rm10^{11.3} - 10^{11.5} \\, L_\\odot$, dust masses of $\\rm 10^8 - 10^9 \\, M_\\odot$ and temperatures from 18 K to 80 K, similar to ULIRGs and (sub)mm detected QSOs. They find that the FIR properties can not be predicted from shorter wavelengths,  because at rest-frame frequencies $0.3 - 20\\, \\mu$m FIR-detected QSOs have similar SEDs of FIR-undetected QSOs at comparable redshifts. Assuming that star formation is powering the FIR emission, about 40\\% of the sample requires $\\rm SFR > 1000 \\, M_\\odot/yr$; yet, for some QSOs the derived SFRs are $\\rm > 5000\\, M_\\odot/yr$, which is unlikely and may imply a non-negligible contribution from the AGN. Multi-components SED fits of a much larger sample of 250 $\\mu$m selected galaxies of HerMES with {\\it Spitzer}/InfraRed Spectrograph spectra show, however, that the presence of a starburst is necessary to account for the total FIR emission, and that the presence of an AGN has no noticable effect on the FIR properties of the host galaxy (Feltre et al. 2013).  Similar studies have been recently applied to 69 QSOs at $z > 5$ by Leipski et al. (2014), who have presented SEDs covering the rest-frame wavelengths from $0.1\\, \\mu$m to $\\sim 80 \\, \\mu$m, including {\\it Spitzer} and {\\it Herschel} observations. At rest-frame wavelengths $\\rm \\ge 50 \\, \\mu m$, the SEDs of the 10  sources which have been detected in at least 4 of the 5 {\\it Herschel} bands (with PACS at 100 and 160 $\\mu$m, and with SPIRE at $\\rm 250, 350, and \\, 500\\,\\mu$m) are completely dominated by a luminous cold dust component with $\\rm 47 \\, K  \\le  T_{FIR} \\le 60 \\, K$ and $\\rm L_{FIR} \\sim 10^{13} \\, L_\\odot$. The origin of the FIR component is further analyzed by stacking the SEDs of {\\it Herschel}-detected (10 objects detected at $\\rm 160, 250, and \\, 350 \\,\\mu m$), partly {\\it Herschel}-detected (14 objects detected at $\\rm 100\\,  and/or \\, 160 \\, \\mu m$) and 33 {\\it Herschel}-non detected sources (33 objects). At rest-frame wavelengths $\\rm \\lambda \\le 15 \\, \\mu m$, the average SEDs of the {\\it Herschel}-detected and partly {\\it Herschel}-detected sub-samples appear very similar,  but partly detected sources show a substantial drop in their average SED at longer wavelengths. The fact that some optical and MIR luminous QSOs do show FIR emission and others do not is interpreted as an indication that star formation is the dominant driver for the additional FIR component observed in {\\it Herschel}-detected QSOs.   All the above studies adopt physically motivated multi-component SED fitting techniques which - however - miss the energy balance between the various components. A notable exception is the work by Dwek, Galliano \\& Jones (2007), who have made a simple decomposition of the observed SED of SDSS J1148 into different components. They used a simple screen model with a Galactic extinction law to calculate the spectrum of the escaping stellar radiation, choosing the magnitude of the extinction so that the total reradiated FIR emission was equal to the total energy absorbed by the dust. In this way they estimated that the inferred FIR luminosity of $\\rm 4.6 \\times 10^{13} L_\\odot$ could be powered by star formation but requires a continuous star formation rate of 2500 $\\rm M_\\odot/yr$ for 400 Myr with a top-heavy Initial Mass Function (IMF).  Here we attempt to make a step forward by using a detailed radiative transfer (RT) model to follow the contribution of the host star formation and the central AGN to dust heating.  Our aim is to to assess to what extent the observed SEDs of $z \\sim 6 $ QSOs in the rest-frame FIR can place constraints on current scenarios for the formation and evolution of the first Super-Massive Black Holes (SMBHs) and their host galaxies.  Information on the star formation rate, the nature of the stellar populations, the mass of gas and dust of the host galaxy are obtained using the data-constrained semi-analytical merger tree model, \\textsc{GAMETE/QSOdust} (Valiante et al. 2011, 2012, 2014). This model allows to follow the star formation history and the mass growth of the nuclear Black Hole (BH) in all the progenitor galaxies that contribute to the hierarchical build-up of $z \\sim 6 $ QSOs. Valiante et al. (2011) show that when the model is applied to SDSS J1148+5251 (hereafter SDSS J1148), one of the best-studied QSOs at $z = 6.4$,  the observed properties of the QSO (BH, gas and dust masses, SFR, stellar mass) are reproduced provided that distinct evolutionary paths are followed. This same conclusion has been recently confirmed to hold for a larger sample of QSOs at $5 < z < 6.4$ (Valiante et al. 2014) which share a common formation scenario: in all the progenitor galaxies of the final host dark matter halo, stars form according to a standard, Salpeter-like IMF via quiescent star formation and efficient merger-driven bursts. At the same time, the central BH grows via gas accretion and mergers with other BHs. When the BH mass reaches a threshold value of $\\rm 2 \\times10^8 - 10^9 \\, M_\\odot$, a strong energy-driven wind starts to clear-up the interstellar medium of dust and gas through a large outflow, damping the star formation rate and un-obscuring the line of sight toward the QSO. Although the estimated outflow rates can be as large as a few $\\rm 1000 \\, M_\\odot/yr$, in agreement with current observational constraints (Maiolino et a. 2012; Cicone et al. 2015), the final star formation rates are comparably large, $\\sim 600 - 2000 \\, \\rm M_\\odot/yr$, due to intense merger-driven bursts. As a result, all the QSOs host galaxies are characterized by final stellar masses in the range $\\rm (4 - 6)\\times 10^{11} M_\\odot$, a factor 3 - 30 larger than the upper limits allowed by the observed dynamical and molecular gas masses.  Although large uncertainties still affect dynamical mass measurements in these high-$z$ galaxies, one possible solution to this so-called {\\it stellar mass crisis} would be to assume that the cold gas is converted into stars at a smaller efficiency than adopted by the scenario described above. In this way, the BH grows faster than the host galaxy and when it reaches a final mass of $\\rm \\sim 10^9 \\, M_\\odot$, the stellar mass is still $\\rm 10^{10} - 10^{11} M_{\\odot}$ (Valiante et al. 2014). Yet, in order to reproduce the observed chemical properties of the host galaxy, in particular the dust mass, the stars must form according to a top-heavy IMF, with a characterstic mass of $\\rm 5 M_\\odot$, that allows to maximize the stellar metal and dust yield (Valiante et al. 2011). More importantly, the final SFRs are $\\rm < 100 \\, M_{\\odot}/yr$, significantly smaller than in the standard model.  It is clear that the two scenarios discussed above result in very different predictions on the relative role of the AGN and star formation as the powering sources of the FIR emission. Hence, any attempt to gain a deeper understanding of the origin of the FIR emission in high-$z$ QSOs may provide important insights on the co-evolution of the first SMBHs and their host galaxies.  The paper is organized as follows: in section \\ref{sec:model} we give a short description of the main features of the semi-analytical model  \\textsc{GAMETE/QSOdust} and we present the standard scenario for SDSS J1148 at $z= 6.4$, using this QSO as a prototype for the general class of high-$z$ QSOs; in section~\\ref{sec:sed} we describe the numerical set-up and the radiative transfer model; in sections ~\\ref{sec:central} - \\ref{sec:dust}, we present the adopted model for the spectrum emitted by the central AGN and by the stars, and the assumed spatial distribution of the dust in the host galaxy; in section \\ref{sec:results} we illustrate and critically discuss our main results and in section \\ref{sec:conclusions} we draw our conclusions.  For consistency with previous works (Valiante et al. 2011, 2014), we adopt a Lambda Cold Dark Matter ($\\Lambda$CDM) cosmology  with $\\Omega_m = 0.24$, $\\Omega_\\Lambda = 0.76$, $\\Omega_b = 0.04$, and $H_0 = 73$~km/s/Mpc. The age of the Universe at a redshift $z = 6.4$ is 900 Myr and the model and observations are scaled using the luminosity distance at this redshift, 64.5 Gpc. ",
        "conclusions": "\\label{sec:conclusions}  One of the main motivations of the present study was to explore to what extent the far infrared properties of high-$z$ QSOs can help constrain current scenarios for the coevolution of the BH and its host galaxy in the first Gyr of cosmic evolution.  In FIR luminous high-$z$ QSOs, such as the $z = 6.4$ QSO SDSS J1148 that we have selected as a representative case, the observed FIR emission is generally attributed to an ongoing strong starburst in the host galaxy, with inferred star formation rates that can range from $\\sim 1000 \\,\\rm M_\\odot/yr$ up to $\\sim 3700\\, \\rm M_\\odot/yr$ (see Table 1 in Valiante et al. 2014). Yet, these figures can only provide a loose upper limit on the star formation rate. In fact, by means of a detailed RT calculation, we show that:  \\begin{itemize} \\item the radiation emitted by the central source, which dominates the observed SED from the UV/optical to near and mid-infrared wavelengths, provides an important source of heating for the dust distributed in the host galaxy, powering at least 30 and up to 70\\% of the observed far infrared emission.  \\item the FIR SED of SDSS J1148 can only be reproduced by models where the star formation rate in the host galaxy at $z = 6.4$ is as large as $\\rm SFR \\sim 900 \\, M_\\odot/yr$, as predicted by the selected standard IMF model (Valiante et al. 2011).  In fact, a top-heavy IMF model, with reduced star formation efficiencies, predicts a SFR at $z = 6.4$ of only $\\rm 66 \\pm 59 \\, M_\\odot/yr$, too small to supply the required heating, in addition to that of the central source, and to power the observed FIR emission.  \\end{itemize}  Hence, we conclude that our analysis supports a scenario for the co-evolution of the first SMBHs and their host galaxies where, during their hierarchical assembly, the evolutionary tracks followed by the systems in the $\\rm M_{BH} - M_{star}$ plane approach the local relation from the bottom (see Fig.~5 in Valiante et al. 2014): during the first period of evolution, at redshifts $8-10 < z < 30$, in all the progenitor galaxies stars form according to a standard, Salpeter-like IMF via quiescent star formation and merger-driven bursts. At the same time, the nuclear BHs grow through mergers with other BHs and gas accretion; hence BHs and their host galaxies grow hand-in-hand, in a symbotic way (Volonteri 2012).  At $5 - 6.4 \\le z \\le 8 - 10$, however, the BH starts growing faster than the stellar bulge and the predicted $\\rm M_{BH} - M_{star}$ evolution progressively steepens, falling within the observed scatter of the local relation. This scenario implies that current dynamical mass measurements may have missed an important fraction of the host galaxy stellar mass.  It is important to note that the transition between the starburst-dominated stage and the QSO-dominated stage of the evolution is regulated by the effects of BH feedback: for the QSOs that we have investigated, we find that when the nuclear BH masses have grown to $\\rm \\sim 2 \\times 10^8 - 10^9 M_\\odot$, strong AGN-driven outflow rates of $\\rm (4 - 6) \\times 10^3 M_\\odot/yr$ damp the star formation rate and un-obscure the line of sight towards the QSO. At these large outflow rates, the host galaxies would be completely depleted of their gas content in less than 20 Myr, shutting down both the star formation rate and the BH activity; as a result, we predict the active and bright QSO phase to last $\\rm \\sim 10^7 yr$ (Valiante et al. 2014). During this relative short phase, the FIR luminosity is expected to respond to the rapidly changing physical conditions in the host galaxy. Our analysis suggest that FIR-luminous QSOs are likely tracing a phase of the evolution where the gas content in the host galaxy can still sustain strong starbursts, with $\\rm 100 M_\\odot/yr < SFR < 1000 M_\\odot/yr$. When the conditions in the ISM are such that the star formation rate fall below $\\rm < 100 M_\\odot/yr$, the FIR emission can only be powered by dust heating from the central source, and we expect a stronger correlation between the UV/optical, NIR/MIR and the residual FIR luminosities.  It is interesting to note that this results resonate with the recent analysis by Leipski et al. (2013, 2014), where the authors find that the average SED of QSOs $z > 5$ with {\\it Herschel} detections, including SDSS J1148, are preferentially found at the high luminosity end of their sample (for $\\rm L_{UV/opt}$ and in particular for $\\rm L_{NIR}$). Moreover, as discussed in section \\ref{sec:intro}, they also find that some optical and MIR luminous QSOs do show FIR emission and others do not. This is an indication that star formation is the dominant driver for the additional FIR component observed in their {\\it Herschel}-detected sources. Following a similar approach, Barnett et al. (2015) has recently analyzed the SED of QSO ULAS J1120+0641 at  $z = 7.1$. This object has been detected by {\\it Herschel} PACS at $100$ and $160 \\, \\mu$m but not by {\\it Herschel} SPIRE at $250$, $350$ and $500 \\, \\mu$m, hence it would be a member of the partly ({\\it Herschel}) detected objects, according to the classification by Leipski et al. (2014). On average, these sources are optically luminous AGN with strong NIR and MIR emission, but do not show a bright FIR emission, consistent with the idea that the central source is likely contributing significantly, if not dominantly, to the FIR emission. Using local scaling relations, Barnett et al. (2015) estimate a star formation rate in the range $\\rm 60 - 270 \\, M_\\odot/yr$ from the [CII] line luminosity and the $\\rm 158\\, \\mu$m continuum luminosity and conclude that, at the time of observation, the BH was growing in mass 100 times faster than the stellar bulge. This is consistent with the idea that the source was likely caught in action when the QSO-driven wind has already significantly depleted the host galaxy of its gas content, suppressing star formation.  Our results are in good agreement with another study done on the same object by Li et al. (2008). By applying a 3D radiative transfer code to a hydrodynamical simulation, these authors show that the observed SED and inferred dust properties of SDSS J1148 request vigorous star formation in the merging progenitors and a substantial contribution to dust heating and the FIR luminosity from the AGN. They also suggest that the quasar host should have already built up a large stellar population by $z \\sim 6$, in agreement with our best-fit evolutionary scenario described above. Our study, which is based on a semi-analytical code, complements these previous findings allowing an exploration of a larger parameter space (in terms of stellar IMF and star formation history) and a more sophisticated treatment of the chemical enrichment of the ISM by dust and metals."
    },
    "1402/1402.0756_arXiv.txt": {
        "abstract": "Systems of self-gravitating fermions constitute a topic of great interest in astrophysics, due to the wide field of applications. In this paper, we consider the gravitational equilibrium of spherically symmetric Newtonian models of collisionless semidegenerate fermions. We construct numerical solutions by taking into account the effects of the anisotropy in the distribution function and considering the prevalence of tangential velocity. In this way, our models generalize the solutions obtained for isotropic Fermi-Dirac statistics. We also extend the analysis to equilibrium configurations in the classical regime and in the fully degenerate limit, recovering, for different levels of anisotropy, hollow equilibrium configurations obtained in the Maxwellian regime. Moreover, in the limit of full degeneracy, we find a direct expression relating the anisotropy with the mass of the particles composing the system. ",
        "introduction": "The first evidence of dark matter was found by Zwicky \\cite{zwicky33}, who noted (in studying the Coma cluster of galaxies) that the dispersion in the radial velocity of the galaxies was very large (around 1000 km s$^{-1}$) and that the stars and gas visible within the galaxies did not provide enough gravitational attraction to keep bound the cluster. In order to maintain the galaxies in the Coma cluster Zwicky concluded that the cluster had to contain a large amount of ``dunkle materie,\" i.e., ``dark matter.\" A confirmation of Zwicky's suggestion was furnished about 40 years later from optical and 21-cm observations of rotation curves of spiral galaxies that did not show the Keplerian falloff, implying the presence of an additional mass component and suggesting that spiral galaxies have a massive dark halo extending to several times the radius of the luminous matter and containing most of the total of their mass \\cite{1987bt_book}.  Then, from the rotation curve it is possible to parametrize the distribution of matter within the halo. As a first approximation, we may schematize the halo as a spherical one and neglect the gravitational contribute of the disk. If we furthermore assume a rotation curve with constant velocity $v_{c}$, the corresponding dark halo density profile corresponds to $\\rho(r) \\sim 1/r^{2}$. The \\emph{N}-body simulations are a powerful tool to study the behavior of particles forming dark halos and give us a better parametrization of the halo density, that we may write as \\begin{equation} \\label{n_body} \\rho(r) = \\frac{\\rho_{0}}{{(\\frac{r}{r_{s}})}^\\lambda {[1 + {(\\frac{r}{r_{s}})}^\\mu]}^\\nu}\\ , \\end{equation} where $\\lambda$, $\\mu$ and $\\nu$ are parameters defining the profile, while $r_{s}$ is the radius scale (the scale in which the profile changes shape, connected with the core radius). For $\\lambda$ = 1, $\\mu$ = 1, and $\\nu$ = 3, Eq.(\\ref{n_body}) corresponds to the Hernquist model \\cite{1990hernquist} while the model proposed by Navarro, Frenk and White \\cite{nfw96,nfw97} has $\\lambda$ = 1, $\\mu$ = 1, and $\\nu$ = 2 and Moore \\emph{et al.} \\cite{mgqsl98,mgqsl99} proposed the model with $\\lambda$ = 1.5, $\\mu$ = 1, and $\\nu$ = 1. These proposals have an unpleasant feature, due to the presence of a cusp when the density is evaluated at the center: to solve this problem, Burkert \\cite{burkert_profile} has proposed a different expression for the density profile of dark matter \\begin{equation} \\label{burkert} \\rho(r) = \\frac{\\rho_{0} r_{0}^{3}}{(r + r_{0}) (r^{2} + r_{0}^{2})}\\ , \\end{equation} where the parameters $\\rho_{0}$ and $r_{0}$ are free and represent, respectively, the central density and the radius scale, in the same sense of Eq.(\\ref{n_body}). Equation (\\ref{burkert}) is similar to the isothermal profile in the limit $r \\ll r_{0}$ (\\cite{burkert_profile} and references therein) and predicts a finite central density.  In order to find candidates for particles forming halos, many hypotheses have been advanced. In the field of cosmology \\cite{2009rich_book}, the theory of structure formation assumes that dark matter is composed by massive particles weakly interacting with matter, named WIMPs (predicted, moreover, by supersymmetric extensions of the standard model of particle physics). Anyway, there are other candidates to explain the formation of dark matter, like axions \\cite{1988cheng} - invented to prevent \\emph{CP} violation in the strong interactions, and massive compact Halo objects (MACHOs), that represent the simplest way to hide baryons and detectable from their gravitational lensing \\cite{1986pacz} and cold gas \\cite{1997cpf}, an alternative to MACHOs (in this case, baryons are hidden in small clouds comprised of primordial helium and molecular hydrogen).  There is indeed another way to study the properties of dark halos, by considering a different approach which takes into account statistical mechanics. A particular interest is due to systems made by massive neutrinos, especially in connection with the problem of the formation and stability of large scale structures like galaxies and clusters of galaxies \\cite{cmcc72,cmcc73}. This case represents the starting point of the discussion of dark matter in terms of fermions. If we generally think about a system of fermions (assumed to be collisionless), it is not possible to use Fermi-Dirac distribution without considering a cutoff in kinetic energy, because masses and radii of definite configurations are not finite. For this reason arises the necessity to introduce a cutoff energy in Fermi-Dirac statistics \\cite{1983rufs}: \\begin{equation} \\label{fd_cutoff} \\begin{cases} f=\\frac{g}{h^{3}} \\frac{1 - e^{(\\epsilon - \\epsilon_{c})/kT}}{e^{(\\epsilon - \\mu)/kT} + 1} & \\text{for $\\epsilon\\leq\\epsilon_{c}$} \\\\ f=0 & \\text{for $\\epsilon > \\epsilon_{c}$}\\ , \\end{cases} \\end{equation} where $\\epsilon$ is the kinetic energy of a single particle and $\\epsilon_{c}$ the cutoff energy, $g = 2s + 1$ is the multiplicity of quantum states, $\\mu$ is the chemical potential and $T$ is the temperature; in the limit $\\epsilon_c\\rightarrow\\infty$ we recover the semidegenerate Fermi-Dirac distribution. In the fully degenerate case, we have a natural cutoff velocity associated to the Fermi energy. Then, the question of the relation between the Fermi energy $\\epsilon_F$ and the cutoff energy $\\epsilon_c$ is arising in order to better understand how the particles are distributed when $T\\rightarrow 0$. We know that the energy variation due to a single particle (at constant entropy $S$ and volume $V$) cannot be larger than the maximum energy that a particle can have. This condition can be expressed by the inequality $(\\partial U/\\partial N)_{S,V}\\leq\\epsilon_c$, with $\\epsilon_c$ being the maximum energy for a single particle. By the definition of the chemical potential $\\mu = (\\partial U/\\partial N)_{S,V}$ we get the important relation $\\mu\\leq\\epsilon_{c}$, valid at every regime and trivially verified in the nonquantum limit where $\\mu<0$. In the fully degenerate case, indeed, from the condition $\\epsilon_F =\\mu$ we obtain that $\\epsilon_F\\leq\\epsilon_c$: namely the cutoff energy never affects the value of the Fermi energy. Moreover, in the interval between $\\epsilon_F$ and $\\epsilon_c$ the density in phase space is zero, leading to the conclusion, valid only in the fully degenerate case, that the cutoff energy must coincide with the Fermi one. Therefore, the relation must be $\\epsilon_F =\\epsilon_c$.  A system described by Eq.(\\ref{fd_cutoff}) is characterized by an isotropic distribution of velocities of particles. Nevertheless, the presence of anisotropies in the momentum distribution, depending on the angular momentum (of particles forming the halos) $L = mrv_{t}$ ($m$, $r$, and $v_{t}$ are,  respectively, mass, radius, and tangential velocity), was found by \\emph{N}-body simulations \\cite{nav_2010} and observational data \\cite{hhpmekv}. However, anisotropy is not a novelty in astrophysics. Michie \\cite{michie61,michie63}, for example, studied the consequences of anisotropy for what concerns stellar systems, and, successively, many authors advanced proposals in this sense by considering stellar systems described by an anisotropic Boltzmann-like distribution function (\\cite{bkmv09} and references therein). Anisotropic models had been also advanced for stellar models, for neutron stars (e.g., \\cite{hh75}) or by considering generalized polytropic models (e.g., \\cite{herr_barreto13a}). More generally, Lingam and Nguyen \\cite{linguyen13} found a method consisting of two different approaches for deriving anisotropic distribution functions and they applied it to Veltmann, Hernquist, and Plummer models. Furthermore, a discussion about the effects of the presence of a local anisotropy in compact objects is in Ref. \\cite{herrera_sant97}. Extended models, generalizing the results of \\emph{N}-body simulations [see Eq.(\\ref{n_body})] with the presence of constant anisotropy, can be found in Ref. \\cite{ma_he08} and examples of theoretical models, described by the anisotropic version of Eq.(\\ref{fd_cutoff}), are in Refs. \\cite{1989ralston,1992imrs}). It is also important to mention models where the source of anisotropy is represented by the action of external fields \\cite{canchiu_68a,canchiu_68b,2010fikps,2012dms}.  In this paper we study the gravitational equilibrium of static Newtonian configurations with spherical symmetry, composed by collisionless semidegenerate Fermi gas, in the presence of a cutoff term in the distribution function. We analyze different levels of anisotropies in the momentum phase space, addressing the relativistic solutions of anisotropic equilibrium models in a forthcoming paper. Finally, for the sake of completeness, it is important to stress the possibility of considering classical distribution functions, due to large values of the predicted mass of the particles \\cite{deV_sanchez}. This question has been considered in the Appendix. ",
        "conclusions": "We have constructed models of anisotropic Newtonian ensembles of collisionless semidegenerate fermions with a distribution function which generalizes the Fermi-Dirac distribution with the energy cutoff, to describe the density profiles of equilibrium configurations which may be serious candidates for describing galactic dark halos. The distribution of dark matter in galactic halos may be estimated by using a method introduced by Persic and Salucci \\cite{ps88}, i.e., the mass decomposition from rotational curves of galaxies. In fact, there is a clear connection between the disk-to-halo mass ratio, $M_{disk}/M_{halo}$, (at optical radii) and the luminosity $L_{B}$ being $M_{disk}/M_{halo} \\propto L_{B}^{2/3}$ \\cite{ps90a,ps90b}. This method therefore relates photometric measurements to distribution of matter.  The results displayed in Tabs.\\,I$-$V and in the figures indicate clearly that, by varying values of \\emph{a}, the anisotropy in the distribution of velocities of particles influences in a different way the behavior of the equilibrium configurations constructed. For high values of \\emph{a} (1, 0.5), we note that the motion of particles is not influenced by anisotropy in a considerable way. The $\\eta$ parameter unlikely reaches its minimum value even if, for \\emph{a} = 0.5, in the intermediate zones of configurations ($0.4<r/R<0.8$), the descent arrives also to values 0.75 $-$ 0.8 (this is due to a decrease of the level of degeneracy within the ensemble). We can further note that the two values 1 and 0.5 do not affect in a considerable way the spatial extension of the particles and the total mass in the whole. From the tables we note indeed that the difference in terms of radii is very small (the discard is lower than 4$\\%$) and still lower in the case of mass (we often note some values are practically the same). About density profiles, we recover the result of Ruffini and Stella \\cite{1983rufs}. In the classical limit ($\\theta_{0} \\ll 0$) they tend to the typical behavior of the isothermal sphere and indeed the sizes of the ensembles of particles increase in a considerable way.  The situation becomes more interesting when \\emph{a} = $10^{-1}$ because the anisotropy is more evident. From the figures we may note that the $\\eta$ parameter approaches more its minimum value and for a larger spatial extension. The density profiles change drastically form, by evidencing the presence of hollow regions \\cite{1991ralsmith,nguyen_ped13} and indicating that \\emph{a} = $10^{-1}$ can make up a critical value for the trigger or not of ``hollowness.\" The results of the numerical integration show also the typical sizes and the mass of the configurations are smaller than the isotropic case, and a higher degree of anisotropy in the momentum distribution tends to ``condense\" spatially the particles.  By still decreasing the value of the anisotropy parameter (\\emph{a} = $10^{-3}$) we note as the hollowness of the density profiles is decidedly more evident. The maximum of the function $\\rho/\\rho_{0}$ increases in absolute value and its position is projected towards the peripheric zones of the ensembles of particles at decreasing the level of degeneracy. The motion of particles suffers more of the presence of the anisotropy. The two isotropic regions (central and peripheric) tend to become thinner and the decrease of $\\eta$ from the maximum value to minimum is sharper than the preceding case. From data recorded in the tables we may see that a decrease by a factor of 100 of the anisotropy parameter implies a reduction of the radius and the mass by a factor about of 10, justifying the general increase of the density of the system and, in particular, of its maximum value. When \\emph{a} = $10^{-5}$, the behavior of the density profiles and the $\\eta$ parameter is more accentuated than the preceding case: the motion of particles becomes totally anisotropic and the density of the system still increases.  In the limit of full degeneracy (see Sec. 5), we have derived an expression in which the dependence of the mass by the anisotropy in distribution of momentum is explicit. In the limit $a \\rightarrow 0$ and for values of matter density comprised between ($\\sim 10^{-29} - \\sim 10^{-17}$) g/cm$^{3}$ (with $W_{0} = 50$) we obtain $m \\geq 0.05$ eV and $m \\geq 48.6$ eV while, in the limit of complete isotropy ($a \\geq 0.5$), for the same values of $\\rho$ and $W_{0}$ we obtain $m \\geq 0.6$ eV and $m \\geq 616$ eV. The first three values of lower limit agree with the hypothesis that dark matter (both for galactic halos and for cosmic structures) is made by the lightest particles, like gravitinos and neutrinos, while the last value indicated is close to values proposed for more massive candidates for dark matter (e.g. axinos).  The level of degeneracy affects, as well as the degree of anisotropy, the behavior of the density profiles and we may summarize this as follows: from the degenerate systems (large values of $W_{0}$ and $\\theta_{0}$) to semidegenerate ones ($\\theta_{0} \\sim 0$), we observe a shifting of the maximum of the density function towards the external regions of the equilibrium configuration, by noting an increase of the value of the maximum. Finally, for large and negative values of $\\theta_{0}$ (classical limit), we recover the same result obtained by BKMV09.  \\appendix"
    },
    "1402/1402.0083_arXiv.txt": {
        "abstract": "While the shape of the extinction curve on the infrared is considered to be set and the extinction ratios between infrared bands are usually taken to be approximately constant, a recent number of studies point either to a spatially variable behavior on the exponent of the power law or to a different extinction law altogether. In this paper, we propose a method to analyze the overall behavior of the interstellar extinction by means of the red-clump population, and we apply it to those areas of the Milky Way where the presence of interstellar matter is heavily felt: areas located in $5^{\\circ}<l<30^\\circ$ and $b=0^\\circ$. We show that the extinction ratios traditionally used for the near infrared could be inappropriate for the inner Galaxy and we analyze the behavior of the extinction law from $1\\mu m$ to 8$\\mu m$. ",
        "introduction": "The extinction law is a useful analytical formula that expresses the amount of light from a given source, in magnitude units, absorbed by the interstellar medium as a function of wavelength. In the infrared, this function $A_{\\lambda}$ is usually modeled as a simple power law of the form\\footnote{See Section \\ref{elaw} for an explanation of the choice of exponents} $A_{\\lambda}\\propto\\lambda^{-\\beta}$ \\citep[see for example][]{RL85,C89}, although other more complex relations have also been proposed \\citep{FM09}.  When the bandpasses under consideration are wide (as in most photometric systems, often several hundred angstroms), $A_{\\lambda}$ is often computed at a given wavelength representative of each filter. The calculation of this value is problematic: it should be the isophotal wavelength \\citep[as defined, for example, in ][]{gol74}, a value that depends on the considered spectral energy distribution. This implies that for every spectral type there is a different isophotal wavelength, and normally an average over several types or the value for a given one (normally a solar-like star) are used instead, or simply the central bandpass' wavelength, much more easily calculated but less representative of the behavior of the extinction under said filter.  Ideally, the extinction law should be tailored for every photometric system, but usually a standard relation \\citep[such as the one proposed by][]{RL85} is used instead. This relation is normally parameterized through several $A_\\lambda/A_\\mathrm{V}$ ratios, relating the extinction for a given band to that of the visible $V$ filter. In this paper, we propose a method for the calculation of these coefficients, assuming a $(\\lambda,A_{\\lambda})$ relation, through the use of red clump giants (RCGs hereafter), and we apply it to 2MASS data \\citep{SK06} in the inner Galactic plane ($5^{\\circ}<l<30^{\\circ}, |b|\\leq0.25^{\\circ}$), where extinction is more severe. Although 2MASS has coverage of the whole disk, this range complements that of \\citet{SH09}, avoiding the innermost Galaxy where source density severely affects 2MASS' depth. Further on, these results will be extended to the mid-infrared using GLIMPSE \\citep{B03} data. ",
        "conclusions": "In this work we study the large scale behavior of the infrared extinction on the inner Milky Way. We present a method to obtain extinction ratios that does not rely on the choice of wavelength for a given set of filters, and applying it to 2MASS and GLIMPSE data, we find that:  \\begin{itemize} \\item[-]The interstellar matter seems to be more transparent in the infrared than previously thought, and the traditional visible to infrared ratios of \\citet{RL85} and \\citet{C89} do not provide a good representation of its evolution over wavelength. \\item[-]Using this method, there is no evidence for an azimuthal variation of the extinction law toward the innermost Galaxy ($0^\\circ<l<30^\\circ$). These variations are present either on a spatial scale that is too small (hence smeared out in our data) or too large (when comparing the inner and outer disk, for example). \\item[-]This change happens in $R_{GC}$, as near the solar system the color excess ratios are quite similar to those of \\cite{RL85}, while towards the Galactic center these change to higher values. Although a physical explanation for this is beyond the scope of this paper, an obvious hypothesis is to tie this change in behavior with the first cross of a spiral arm, where high density molecular regions are likely to affect light in a different way than local, diffuse interstellar medium. \\item[-]Neither a single exponent for a power law nor the expression from \\cite{FM09} seem to be able to account for the overall shape of $A_{\\lambda}$ between 1 and 8$\\mu m$, although a proper calibration of the intrinsic colors for the  red clump stars is needed to further assess this claim. \\end{itemize}  All of this has profound implications on what we know about the structure of our Galaxy, particularly regarding its large-scale geometry, as much of the studies on the bulge and bar assume a much smaller $A_{J}/A_{K}$ ratio than the one found here. This implies that the distances estimated using this ratio will be underestimated if compared when compared to the values that could be derived with the $A_{J}/A_{K}$ from this paper. It can be easily shown that: \\begin{equation} \\frac{d}{d_{\\mathrm{RL}}}\\sim10^{0.34\\cdot E(J-K_\\mathrm{S})} \\end{equation} Where $d_{\\mathrm{RL}}$ is the distance derived with the extinction law from \\cite{RL85} and $d$ the same value but using the results from this paper. For very reddened fields, this ratio can be as high as 2. Yet there is an extra element that we should take into account: as can be seen in Figure \\ref{ratall}, there is evidence that the color excess ratio (and also $A_{J}/A_{K}$) is not constant along the line of sight and that for the first few kiloparsecs from the Sun, values from \\cite{RL85} work well. Being so, we cannot do a straight comparison between distances along the line of sight derived using the classical extinction law and the one presented here, as there is only a fraction of the light path where they differ."
    },
    "1402/1402.1085_arXiv.txt": {
        "abstract": "{exoplanets; atmospheric models; space missions} Planetary science beyond the boundaries of our Solar System is today in its infancy. Until a couple of decades ago, the detailed investigation of the planetary properties was restricted to objects orbiting inside the Kuiper Belt. Today we cannot ignore that the number of known planets has increased by two orders of magnitude, nor that these planets resemble anything but the  objects present in our own Solar System. Whether this fact is the result of a selection bias induced by the kind of techniques used to discover new planets -- mainly radial velocity and transit -- or simply the proof that the Solar System is a rarity in the Milky Way, we do not know yet. What is clear, though, is that the Solar System has failed to be \\textit{the} paradigm  not only in our Galaxy but even ``just'' in the solar neighbourhood. This finding, although unsettling, forces us to reconsider our knowledge of planets under a different light and perhaps question  a few of the theoretical pillars on which we base our current ``understanding''.   The next decade will be critical to advance in  what we should perhaps call \\emph{Galactic planetary science}. In this paper, we review highlights and pitfalls of our current knowledge of this topic and elaborate on how this knowledge might arguably evolve in the next decade. More critically, we identify what should be the mandatory scientific and technical steps to be taken in this fascinating journey of remote exploration of planets in our Galaxy. ",
        "introduction": "If we had to select a single word to define the field of exoplanets, that word would be \\emph{revolutionary}. During the past years, over one thousand planets have been found around every type of stars from A to M, including pulsars and binaries. Being the leftover of the stellar formation processes, planets appear to be rather ubiquitous and, in reality, the presence of a host star is not even a mandatory circumstance. The current statistical estimates indicate that, on average, every star in our Galaxy hosts at least one planetary companion \\citep{Cassan2012}, i.e.  our Milky Way is crowded with one hundred billion planets.  The most revolutionary aspect of this young field is the discovery that the Solar System does not appear to be the paradigm in our Galaxy, but rather one of the many possible configurations we are seeing out there. These include planets completing a revolution in less than one day, as well as planets orbiting two stars or moving on trajectories so eccentric to resemble comets. This variety of stellar and orbital parameters convert into planetary temperatures which span over two orders of magnitudes. Unexpectedly, planetary sizes and masses do not appear to be ``quantised\" as it happens in the Solar System, where the terrestrial planets are well separated from Neptune and Uranus and those are, in turn, well distinct from Jupiter and Saturn. Instead, a continuum of sizes and masses appear to exist, from the super-Jupiter down to the sub-Earth objects \\citep{Batalha2013, Fressin2013}.  While the relative frequency of  ``odd'' planets compared to the ``normal'' ones -- assuming the Solar System planets represent the normality -- might be the result of some selection effects caused by the detection techniques used so far -- mainly radial velocity and transit -- it is undoubtable that a great diversity of  planets does exist around other stars. In the short term we should be able to shed light on this issue. The European Space Agency's GAIA mission is expected to find several thousands new planets through astrometry \\citep{Sozzetti2011, Sozzetti2013}, a technique sensitive to planets lying in a different region of the parameter space compared to transit and radial velocity, in particular to planets at intermediate separation -- typically a few AUs -- from their mother star. The instruments ESO-VLT SPHERE \\citep{Beuzit2008}, Gemini Planet Imager \\citep{Hartung2013} and SUBARU SCExAO \\citep{Jovanovic2013} were built to detect young, massive planets at large separation from the stars, a regime not yet well explored till now.   With these numbers and premises, emphasis in the field of exo-planetary science must shift from discovery to understanding: understanding the nature of exo-planetary bodies, and understanding their history. The following fundamental questions need to be addressed: \\begin{itemize} \\item\tWhat is the origin of the observed exoplanet diversity? \\item\tHow and where do exoplanets form? \\item\tWhat are the physical processes  responsible for exoplanet evolution? \\end{itemize} In all disciplines, taxonomy is often the first step toward understanding, yet we do not have to date  even a simple taxonomy of planets and planetary systems. For planets transiting in front of their parent stars -- of which over 400 are known today -- the simplest observables are the planetary radius and, when combined with radial velocity, the mass. Mass and radius allow to estimate the planetary bulk density. From Fig. \\ref{mr} it is evident that even gas giants have a broad range of interior structures and core composition, as shown from the different bulk densities observed \\citep{Guillot2005, Fortney2007}. While this has stimulated very interesting theoretical work on planetary interior and equations of state of hydrogen at high pressure and temperature, the implications on e.g. planetary formation and evolution mechanisms is still unclear. Most likely, the different bulk densities reflect the different nature and size of the planet's core, which in turn will depend on both the formation mechanism and the ``birth distance'' from the parent star. Objects lighter than ten Earth masses (super-Earths/sub-Neptunes, right-hand panel of Fig. \\ref{mr}) are even more enigmatic, as they often can be explained in different ways \\citep{Adams2008, Grasset2009, Valencia2013}. Among those, Kepler-10 b, Kepler-78 b, CoRoT-7 b and 55 Cnc e all have high densities and orbit G-stars like our Sun with periods of less than one day. By contrast, GJ 1214 b, Kepler-11 d, e, f have much lower densities and are subject to less intense insolation because of their longer period or cooler parent star. In the next years, dedicated space missions such as NASA TESS \\citep{Ricker2010} and ESA Cheops \\citep{Cheops} combined with radial velocity surveys, will measure the sizes and masses of a few thousands new planets, completing the current statistics of available planetary densities in the solar neighbourhood down to the terrestrial regime.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.8\\textwidth]{mass-radius} \\caption{\\small  Masses and radii of currently confirmed transiting planets \\citep{Howard2013}. Extrasolar planets are denoted by red circles and Solar System planets are represented by green triangles. Green and brown lines denote Earth-like composition (67\\% rock, 33\\% iron) and Mercury-like composition (40\\% rock, 60\\% iron).  } \\label{mr} \\end{center} \\end{figure}  As explained above, density is a very important parameter, but alone it cannot be used as a discriminant of the variety of cases we are seeing out there. We need additional information to proceed. The other key observable for planets is the chemical composition and state of their atmosphere. Knowing what are atmospheres made of is essential to clarify, for instance, whether a planet was born in the orbit it is observed in or whether it has migrated a long way;  it is also critical to understand the role of stellar radiation on escape processes, chemical evolution and global circulation. To date two methods can be used to sound exoplanetary atmospheres: transit and eclipse spectroscopy and direct imaging spectroscopy. These are very complementary methods and we should pursue both to get a coherent picture of planets outside our solar system (Fig. \\ref{transit}).  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.55\\textwidth]{transit} \\includegraphics[width=0.44\\textwidth,trim=23cm 0cm 0cm 0cm, clip=true]{DI} \\caption{ Left: phase curve of the hot-Jupiter HAT-P7b while orbiting around its mother star as observed by Kepler \\citep{Borucki2009}. The transit and eclipse events are occurring at   10h and 35h respectively. Right: image of the planet $\\beta$-Pic b, located 8 to 15 astronomical units from the star, as observed with instrument VLT-Naco \\citep{Lagrange2010}. }  \\label{transit} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1402/1402.3755_arXiv.txt": {
        "abstract": "We study two types of entropic-force models in a homogeneous, isotropic, spatially flat, matter-dominated universe. The first type is a `$\\Lambda(t)$ type' similar to $\\Lambda(t)$CDM (varying-lambda cold dark matter) models in which both the Friedmann equation and the acceleration equation include an extra driving term. The second type is a `BV type' similar to bulk viscous models in which the acceleration equation includes an extra driving term whereas the Friedmann equation does not. In order to examine the two types systematically, we consider an extended entropic-force model that includes a Hubble parameter ($H$) term and a constant term in entropic-force terms. The $H$ term is derived from a volume entropy whereas the constant term is derived from an entropy proportional to the square of an area entropy. Based on the extended entropic-force model, we examine four models obtained from combining the $H$ and constant terms with the $\\Lambda(t)$ and BV types. The four models agree well with the observed supernova data and describe the background evolution of the late universe properly. However, the evolution of first-order density perturbations is different in each of the four models, especially for low redshift, assuming that an effective sound speed is negligible. The $\\Lambda(t)$ type is found to be consistent with the observed growth rate of clustering, in contrast with the BV type examined in this study. A unified formulation is proposed as well, in order to examine density perturbations of the two types systematically. ",
        "introduction": "$\\Lambda$CDM (lambda cold dark matter) models are the simplest cosmological model which can explain an accelerated expansion of the late universe \\cite{PERL1998ab,Riess1998_2004,Riess2007SN1,WMAP2011,Planck2013}. However, the standard $\\Lambda$CDM model (which assumes a cosmological constant $\\Lambda$ and an additional energy component called `dark energy') suffers from several theoretical problems, e.g., the cosmological constant problem, the cosmic coincidence problem, etc. \\cite{Weinberg1989}. In order to solve the problems, various cosmological models have been suggested, using alternative dark energy, modified gravity, etc. (see, e.g., Refs.\\ \\cite{Weinberg1,Roy1,Miao1,Bamba1,Sola_2013b} and references therein).  In those models, two types of cosmological models have been extensively examined in an effort to explain an accelerated expansion of the universe. The first type is related to $\\Lambda (t)$CDM models which assume a variable cosmological term $\\Lambda(t)$ \\cite{Freese1-Fritzsch1,Overduin1,Sola_2002,Sola_2003,Sola_2004,Sola_2009,Sola_2011a,Sola_2011b,Sola_2013a,Sola_2013b,Sola_2013c,Sola_2014b,Lima_2014a}. In this model, a varying $\\Lambda(t)$ (which corresponds to an extra driving term) is added to both the Friedmann equation and the Friedmann--Lema\\^{i}tre acceleration equation, instead of the cosmological constant $\\Lambda$. We call this the `$\\Lambda(t)$ type' of cosmological model. The second type is related to bulk viscous models (which assume a bulk viscosity of cosmological fluids) \\cite{Davies3,Weinberg0,Murphy1,Barrow11,Barrow12,Lima101,Zimdahl1,Arbab1,Brevik1,Brevik2,Nojiri1,Meng1,Fabris1,Colistete1,Barrow21,Meng2,Avelino1,Hipo1,Avelino2,Piattella1,Meng3,Pujolas1,Odintsov1,Odintsov2,Odintsov3,Odintsov4} and CCDM models (which assume a creation of cold dark matter)  \\cite{Lima-Others1996-2008,Lima2010,Lima2010b,Lima2011,Lima2012}. In the bulk viscous and CCDM models, the Friedmann equation does not include an extra driving term because dissipation processes are assumed. We call this the `BV (bulk viscous) type'.  Recently, Easson \\textit{et al.} have proposed an entropic-force model as an alternative explanation for the accelerated expansion of the universe \\cite{Easson1,Easson2}. We expect that the entropic-force model \\cite{Easson1,Easson2,Koivisto1,Koma4,Koma4bc,Koma5,Cai1_Cai2_Qiu1,Casadio1-Costa1,Basilakos1,Lepe1,Sola_2014a} is related to $\\Lambda(t)$ and BV types. In the entropic-force model, an extra driving term, i.e., an entropic-force term, is derived from the usually neglected surface terms on the horizon of the universe \\cite{Easson1}. The entropic-force term can explain the accelerated expansion, without introducing new fields and an exotic energy component of the universe such as dark energy. Instead of dark energy, the entropic-force model assumes that the horizon of the universe has an entropy and a temperature due to the information holographically stored there \\cite{Easson1}. For example, the Bekenstein entropy (area entropy) and the Hawking temperature are used in the original entropic-force model \\cite{Easson1}. The obtained entropic-force term is usually added to both the Friedmann and acceleration equations \\cite{Easson1,Easson2,Koivisto1}. (Note that the entropic-force considered here is different from the idea that gravity itself is an entropic-force \\cite{Padma1,Verlinde1}.) Accordingly, the original entropic-force model is $\\Lambda(t)$ type. The present authors have proposed a modified entropic-force model \\cite{Koma5} assuming a generalized black-hole entropy proportional to its volume, based on appropriate nonadditive generalizations \\cite{Tsallis2012}. The obtained entropic-force terms behave as if they were an extra driving term for bulk viscous models. Therefore, the present authors have assumed dissipation processes similar to the bulk viscous model \\cite{Koma5}. That is, the modified entropic-force model corresponds to the BV type. In this way, entropic-force models proposed so far can be categorized into $\\Lambda(t)$ and BV types, using different formulations.  In entropic-force models, entropic-force terms depend on the class of entropy. For example, $H^{2}$ terms are derived from an area entropy \\cite{Easson1}, whereas $H$ terms are derived from a volume entropy \\cite{Koma5}, where $H$ is the Hubble parameter. The entropic-force terms affect the background evolution of the universe. Note that we do not discuss inflation of the early universe. In fact, entropic-force models which include $H^{2}$ terms cannot describe a decelerating and accelerating universe \\cite{Basilakos1,Sola_2013a,Koma4}. On the other hand, a modified entropic-force model which includes $H$ terms can describe a decelerating and accelerating universe \\cite{Koma5}. Of course, the entropic-force terms affect density perturbations as well. For example, Basilakos \\textit{et al.} have recently shown that the original entropic-force model (which includes $H^{2}$ terms) does not describe cosmological fluctuations properly without the inclusion of a constant term \\cite{Basilakos1}. Also, they have found that $\\Lambda(t)$CDM models (similar to the original entropic-force model) are not consistent with the structure formation data \\cite{Sola_2009}. Furthermore, Li and Barrow have explained that bulk viscous models (which include $H$ terms) are difficult to reconcile with astronomical observations of structure formations \\cite{Barrow21}. The previous works suggest that it is necessary to consider not only an $H$ term but also a constant entropic-force term. (Entropic-force models which include $H$ terms have been recently investigated by Basilakos and Sol\\`{a} \\cite{Sola_2014a}.)  The constant term has been examined in $\\Lambda(t)$CDM \\cite{Sola_2009,Sola_2011b,Sola_2013a,Sola_2013b,Basilakos1} and CCDM models \\cite{Lima-Others1996-2008,Lima2010,Lima2010b,Lima2011,Lima2012}, whereas $H$ terms have been investigated in $\\Lambda(t)$CDM \\cite{Sola_2009} and bulk viscous models \\cite{Davies3,Weinberg0,Murphy1,Barrow11,Barrow12,Lima101,Zimdahl1,Arbab1,Brevik1,Brevik2,Nojiri1,Meng1,Fabris1,Colistete1,Barrow21,Meng2,Avelino1,Hipo1,Avelino2,Piattella1,Meng3,Pujolas1,Odintsov1,Odintsov2,Odintsov3,Odintsov4}. In those works, the influence of the extra driving terms is focused on, and therefore, the difference between the $\\Lambda(t)$ and BV types has not yet been discussed systematically. This is because the two types are usually categorized into different models, e.g., $\\Lambda(t)$CDM, CCDM, and bulk viscous models. (Background evolutions of the universe in the $\\Lambda (t)$CDM and CCDM models have been discussed in Ref.\\ \\cite{Lima_2014a}.) However, it is possible to examine the two types systematically through entropic-force models. Density perturbations of the two types are expected to be different from each other, because each continuity equation of cosmological fluids is different even if the background evolution of the universe is the same. Therefore, in the present study, we examine the properties of the $\\Lambda(t)$ and BV types of entropic-force models. To this end, we consider an extended entropic-force model which includes $H$ and constant terms. The constant term is derived from an entropy proportional to the square of an area entropy \\cite{Koma5}.    The remainder of the present paper is organized as follows. In Sec.\\ \\ref{General Friedmann equations and entropic-force models}, we present the general Friedmann and continuity equations and discuss entropic-force models. In Sec.\\ \\ref{Entropic-force models}, we present a brief review of two types of standard entropic-force model, i.e., the $\\Lambda (t)$ and BV types. In Sec.\\ \\ref{An extended entropic-force model}, in order to examine the two types systematically, we consider an extended entropic-force model which includes $H$ and constant terms. In Sec.\\ \\ref{Density perturbations}, we briefly review the density perturbations of the $\\Lambda(t)$ and BV types, in the linear approximation. In addition, we discuss a unified formulation, in order to examine the density perturbations of the two types systematically. In Sec.\\ \\ref{The present models}, based on the extended entropic-force model, we propose four models obtained from combining the $H$ and constant terms with the $\\Lambda(t)$ and BV types. In Sec.\\ \\ref{Results}, we examine the evolution of the universe in the four entropic-force models. Finally, in Sec.\\ \\ref{Conclusions}, we present our conclusions. ",
        "conclusions": "\\label{Conclusions}  Entropic-force models are categorized into two types. The first is the $\\Lambda(t)$ type similar to $\\Lambda(t)$CDM models, and the second is the BV type similar to bulk viscous models (and CCDM models). In order to examine the two types systematically, we have considered an extended entropic-force model which includes $H$ and constant $C_{\\rm{cst}}$ terms. In particular, we have focused on the $H$ and $C_{\\rm{cst}}$ terms separately, in a homogeneous, isotropic, and spatially flat matter-dominated universe. The constant entropic-force term considered here is derived from an entropy $S_{r4}$ proportional to $r_{H}^{4}$, assuming $S_{r4}$ as one of the possible entropies.  In the present paper, we have examined four models, the $\\Lambda(t)$-$H$, $\\Lambda(t)$-$C_{\\rm{cst}}$, BV-$H$, and BV-$C_{\\rm{cst}}$ models, which are obtained from combining the $H$ and $C_{\\rm{cst}}$ terms with the $\\Lambda(t)$ and BV types. The four models agree well with observed supernova data points and describe a decelerated and accelerated expansion of the universe. In order to examine first-order density perturbations in the four models, we used two formulations proposed by Basilakos \\textit{et al.} \\cite{Sola_2009} and Jesus \\textit{et al.} \\cite{Lima2011}. The two formulations can be summarized using a neo-Newtonian approach and, therefore, we have proposed a unified formulation which helps to observe the two formulations systematically. In addition, we have extended the formulation, in order to study the BV-$H$ model. Consequently, for large $a/a_{0}$, the perturbation growth factor $\\delta$ for the $\\Lambda(t)$-$H$, BV-$H$, and BV-$C_{\\rm{cst}}$ models decreases, whereas $\\delta$ for $\\Lambda(t)$-$C_{\\rm{cst}}$ does not decrease. Therefore, for low redshift, the growth rate for the $\\Lambda(t)$-$H$, BV-$H$, and BV-$C_{\\rm{cst}}$ models tends to be negative and disagrees with the observed growth rate. It is found that for low redshift, $H$ versions (i.e., the $\\Lambda(t)$-$H$ and BV-$H$ models) are not consistent with structure formations, though the $H$ version describes a decelerating and accelerating universe. In contrast, the growth rate for the $\\Lambda(t)$-$C_{\\rm{cst}}$ model agrees well with the observed growth rate. Interestingly, the BV-$C_{\\rm{cst}}$ model disagrees with the observed growth rate, although its background evolution of the universe is the same as that of the $\\Lambda(t)$-$C_{\\rm{cst}}$ model. (Note that we have assumed $c_{\\rm{eff}}^{2} = 0$ in the present study. The BV-$C_{\\rm{cst}}$ model agrees with the observed growth rate if $c_{\\rm{eff}}^{2}$ is considered to be a free parameter, as examined by Jesus \\textit{et al.} \\cite{Lima2011}.)  It is also found that $\\Lambda(t)$ types (the $\\Lambda(t)$-$H$ and $\\Lambda(t)$-$C_{\\rm{cst}}$ models) are consistent with the observed growth rate, in contrast with BV types (the BV-$H$ and BV-$C_{\\rm{cst}}$ models). Therefore, in entropic-force models, the $\\Lambda(t)$ type is likely suitable for describing density perturbations or structure formations. Through the present study, we have revealed fundamental properties of the two types of entropic-force models systematically. Of course, similar cosmological models have been discussed in bulk viscous models \\cite{Barrow21}, CCDM models \\cite{Lima2011}, and $\\Lambda(t)$CDM models \\cite{Sola_2009}. Our results are related to those cosmological models although the theoretical backgrounds are different.                      \\appendix"
    },
    "1402/1402.6346_arXiv.txt": {
        "abstract": "RadioAstron space-ground VLBI observations of the pulsar \\object{B0950+08}, conducted with the 10-m space radio telescope in conjunction with the Arecibo 300-m telescope and Westerbork Synthesis Radio Telescope at a frequency of 324 MHz, were analyzed in order to investigate plasma inhomogeneities in the direction of this nearby pulsar. The observations were conducted  at a spacecraft distance of 330{,}000 km, resulting in a projected baseline of 220{,}000 km, providing the greatest angular resolution ever achieved at meter wavelengths. Our analysis is based on fundamental behavior of structure and coherence functions. We find that the pulsar shows scintillation on two frequency scales, both much less than the observing frequency; but modulation is less than 100\\%. We infer that the scattering is weak, but a refracting wedge disperses the scintillation pattern. The refraction angle of this ``cosmic prism'' is measured as $\\theta_0=1.1 - 4.4$ mas, with the refraction direction being approximately perpendicular to the observer velocity. We show that the observed parameters of scintillation effects indicate that two plasma layers lie along the line of sight to the pulsar, at distances of $4.4 - 16.4$~pc and $26 - 170$~pc, and traveling in different directions relative to the line of sight. Spectra of turbulence  for the two layers are found to follow a power law with the indices $\\gamma_1 = \\gamma_2 = 3.00\\pm 0.08$, significantly different from the index expected for a Kolmogorov spectrum of turbulence, $\\gamma=11/3$. ",
        "introduction": "\\label{sec:intro}  Small-scale fluctuations in the electron density of the interstellar medium (ISM) scatter radio waves from pulsars. We can study properties of both the ISM and the pulsar magnetosphere by measuring frequency-time characteristics of pulsar scintillation on the baseline of the cosmic interferometer resulting from scattering. Interferometry provides for comparison of scintillation at two different places at the same time.  The goal of the present study is to investigate the spatial distribution of scattering along the line of sight to pulsar \\object{B0950+08}, one of the brightest and nearest pulsars. Its distance, measured by parallax, is $262\\pm 5$ pc \\citep{br}. The observations were conducted at a baseline projection of 220{,}000 km and with a spacecraft distance of 330{,}000 km providing angular resolution of 1 mas -- the maximum ever achieved at meter wavelengths, on the longest baseline yet attained for VLBI observations, using the RadioAstron orbiting antenna.  \\subsection{Interstellar Scattering Material}  Observations suggest the presence of three components of scattering material in our galaxy \\citep{sm1, sm2}. The first is unevenly distributed material in the space between spiral arms (component A). The second corresponds to a cavity of depleted electron density extending as far as 300 pc from the sun in the direction perpendicular to the galactic plane and 50 to 100 pc in the Galactic plane (component B) \\citep{snow, bh}. The power spectra of density fluctuations for components A and B are well described by a Kolmogorov spectrum with an index of 11/3 \\citep{ar, shsm}. The third, component C, is located only about 10 pc from the sun and has an increased level of turbulence. This component is responsible for the intra-day variability of quasars at centimeter wavelengths \\citep{dt, ric}.  As discussed in earlier papers \\citep{sm1, sm2}, the main contributor to the scintillation of the nearby pulsars \\object{J0437-4715} and \\object{B0950+08} is component C. The power spectrum of density fluctuations for these pulsars is flatter than the Kolmogorov spectrum. The index of the power-law is $\\gamma = 3.00 \\pm 0.05$ for \\object{PSR~B0950+08} and $\\gamma = 3.46 \\pm 0.20$ for \\object{PSR~J0437-4715}. As shown by these studies, the level of turbulence in component C is a factor of 20 higher than that in the extended region responsible for scintillation of \\object{PSR~B0809+74} \\citep{sm2}. \\citet{phi} assumed that the scattering material in the direction of B0950+08 was uniformly distributed, and found that the level of turbulence an order of magnitude lower than for any previously measured interstellar line of sight. Their results are also consistent with enhanced scattering in component C, at a distance of only 10 pc.  IDV sources show large-amplitude and rapid variability % caused by scintillation in component C. The time scale of these variations is about the same as it would be for pulsars at centimeter wavelengths. \\citet{Lin} suggested that the partially-ionized surfaces of nearby interstellar clouds are responsible for intraday variability. They note that observed intraday variable sources lie behind such clouds; some where clouds collide. The nearby scattering screen for \\object{PSR~B0950+08} may have the same origin; indeed, its line of sight passes through one of the clouds they identify. The far edge of the Local Bubble may also scatter \\object{PSR~B0950+08}. In this direction, the edge of the Local Bubble lies at a distance of 120 to 160 pc  \\citep{lal}. \\citet{sm2} suggested the existence of strong angular refraction in the direction to \\object{PSR~B0950+08}.   \\subsection{Orbiting Antenna and Observations}  Interferometric observation of pulsars with space and Earth antennas, for baselines of several Earth diameters, provides the possibility of localizing the scattering layers and also allows one to evaluate the influence of refraction on the received emission. \\citet{sh2} presents a preliminary theoretical analysis of this approach. For the first such VLBI observations of the Crab pulsar with ground telescopes see \\citet{kon}. As we discuss in Section\\ \\ref{sec:FresnelScales}, the typical scale of the scattering pattern of a weakly-scattered pulsar is the Fresnel scale, larger than an Earth diameter. Space-Earth interferometry affords the possibility of measuring this scale directly, rather than allowing motions of pulsar, Earth and scattering material to carry it across a single antenna. We make use of both interferometric and time-lag analyses in this paper; the earlier work by \\cite{sm2} used only single-dish observations and provides an interesting comparison.  Because scintillation is a stochastic process, observations must be compared with theoretically-predicted distributions. These distributions may not have analytical forms and can be quite difficult to compute \\citep[e.g.,][]{gw,jo}. Consequently, moments of the distributions, particularly in the form of structure functions, provide measures of the fundamental behavior of structure and coherence functions. Section \\ref{sec:ISMScatteringTheory} describes this approach in the present case.   In this paper, we present results obtained from observations of the pulsar \\object{PSR~B0950+08} carried out on 25 January 2012 at a frequency of 324 MHz with the RadioAstron space radio telescope in conjunction with the Arecibo 300-m telescope (AR) and the Westerbork synthesis array (WB). The RadioAstron project is an international collaborative mission involving a free-flying satellite, Spectr-R, carrying a 10-m space radio telescope (SRT) on an elliptical orbit around the Earth. This space telescope performs radio astronomical observations using very long baseline interferometry (VLBI) techniques in conjunction with ground-based VLBI networks. The orbit of the RadioAstron satellite evolves with time. It has an apogee between 280,000 and 350,000 km, a perigee between 7,000 and 80,000 km, a period of 8 to 9 days, and an initial inclination of $51^{\\circ}$. RadioAstron operates at the standard radio astronomical wavelengths of 1.19 -- 1.63 cm (K-band), 6.2 cm (C-band), 18 cm (L-band), and 92 cm (P-band). Technical parameters of the on-board scientific equipment and measured parameters have been described in two main publications \\citep{ak, kar}. ",
        "conclusions": "We carried out successful RadioAstron space-ground VLBI observations of \\object{PSR~B0950+08} on January 25, 2012, at 92~cm with a spacecraft distance of 330{,}000 km and projected interferometer baseline of 220{,}000 km. These measurements represent the highest angular resolution ever achieved in meter wavelength observations. The qualitative difference of the form of the structure function between long and short baselines, as shown in Figure\\ \\ref{fig.5} and discussed in Section\\ \\ref{sec:comparison_with_theory}, suggests the presence of two scattering plasma layers along the line of sight to the pulsar. From analysis of the time and frequency scales of the scintillation, we find that these are located at distances of 4.4 to 16.4\\ pc, and of 26 to 170\\ pc, as discussed in Section\\ \\ref{sec:scales_and_distances}. The nearby layer dominates the temporal structure of the scintillation, while both the nearby and far layers influence the frequency structure of the scintillation. The velocity of the line of sight in the far and nearby layers, projected onto the direction of refraction by the ``cosmic prism'' are 20~km\\,s$^{-1}$ and $-8$~km\\,s$^{-1}$ correspondingly. After correction for the velocities of the Earth and the pulsar, these are in accord with the velocities typical for interstellar clouds. The cosmic prism is described below and in Sections\\ \\ref{sec:cosmicprism}, \\ref{sec:shifts_in_freq_and_time}, \\ref{sec:scales_and_distances}, and\\ \\ref{sec:asymmetry}.  The distance to the far screen suggests that it may be located at the outer wall of the Local Bubble, which lies at about that distance in the direction of \\object{PSR~B0950+08} \\citep{snow,lal}. The distance of the nearby screen suggests that it lies at the ionized surface of a nearby molecular cloud; indeed, such a screen is seen in the direction of the pulsar \\citep{Lin}.  From analysis of the temporal coherence function on short and long baselines, introduced theoretically in Section\\ \\ref{sec:tempcoherencefunction} and computed for our observations in Section\\ \\ref{ref:sec_coherence_function_measurement}, we studied the spectrum of density fluctuations in the two scattering layers. The spectra of density fluctuations for the two layers were found to follow power laws, with indices $\\gamma_1=\\gamma_2=3.00\\pm 0.08$. These indices differ from the Kolmogorov value of $\\gamma = 11/3$. Note that the Kolmogorov spectrum describes more distant scattering media very well. However, our results suggest that nearby material has a flatter spectrum.  We observe evidence for refraction by an interstellar plasma wedge, or ``cosmic prism''. This refraction results in the observed moderate modulation by scintillation of $m=\\Delta I/I <1$, in combination with narrow fractional scintillation bandwidth $\\Delta \\nu/\\nu_0 < 1$. Usually $m<1$ is characteristic of weak scintillation, whereas $\\Delta \\nu/\\nu_0 < 1$ is characteristic of strong scintillation. However, if the characteristic value of the refraction angle $\\theta_0$ by the cosmic prism is greater than the characteristic value of the diffractive or scattering angle, $\\theta_0 >> \\Theta_{scat}$, then the frequency structure of the scintillation is formed by the frequency dependence of the displacement of the beam path, and these two conditions appear together. We describe this theoretical picture in Sections\\ \\ref{sec:cosmicprism} and\\ \\ref{sec:shifts_in_freq_and_time}, and compute parameters of the inferred prism from our observations in Sections\\ \\ref{sec:scales_and_distances} and\\ \\ref{sec:asymmetry}. For \\object{PSR~B0950+08}, we evaluated the angle of refraction as $\\theta_0 = (1.1 - 4.4)$ mas. The refraction is in a direction nearly perpendicular to the velocity of observer."
    },
    "1402/1402.4519_arXiv.txt": {
        "abstract": "It is increasingly apparent that common merger events play a large role in the evolution of disk galaxies at all cosmic times, from the wet accretion of gas-filled dwarf galaxies during the era of peak star formation, to the collisions between large, dynamically-advanced spiral galaxies and their dry companion satellites, a type of interaction that continues to influence disk structure into the present day.  We also live in a large spiral galaxy currently undergoing a series of impacts from an infalling, disrupting dwarf galaxy.  As next-generation astrometry proposes to place our understanding of the Milky Way spiral structure on a much firmer footing, we analyze high-resolution numerical models of this disk-satellite interaction in order to assess the dynamical response of our home Galaxy to the Sagittarius dwarf impact, and possible implications for experiments hoping to directly detect dark matter passing through the Earth. ",
        "introduction": "We visualize the stars associated with the primary Milky Way model in the left panels of Figure~\\ref{fig:disk} and the same system evolved in isolation for 2.65 Gyr (the time of the longest Sgr orbit simulation) in the middle panels. We note that these initial conditions are quite stable during secular evolution, though a weak bar does begin to emerge after $\\sim 2$ Gyr of evolution.  Each of the modeled Sgr progenitors experiences two disk crossings and approaches a third, corresponding to the three orbital pericenters (as indicated by shaded bands). For both models, the first impact occurs in the outer disk at $R_{GC} \\sim 30$~kpc and is responsible for most of the dark matter mass loss in the satellite but very little stellar mass loss. The second occurs $\\sim0.9$ Gyr later at $R_{GC} \\sim 15$~kpc and is responsible for liberating the majority of the stellar stream debris. The final disk crossing is occurring at the present day, about $\\sim1.8$ Gyr after the first disk crossing. We note that the impact times and radii for our models are quite similar to those presented in past work aimed at reproducing the Sgr stream in great detail, although these past models have not included a dark matter component to the massive Sgr galaxy \\citep[see, {\\em e.g.}][]{law_majewski10}.  The left panel of Figure~\\ref{fig:spiral} visualizes the global disk response to the Sgr interaction in the {\\em Light Sgr} model by using a face-on disk rendering color-coded by ratio of the local density to the radial/vertical average at each point. Under-densities are white in this depiction, while over-densities are colored.  The images reveal a prominent bar as well as a complicated array of spirality that reveals itself in the nearby anti-Center as overdensities rotating coherently with the Milky Way disk. The right panel depicts the local density as determined within a 5-degree solid angle cone along the plane as a function of radius from the solar position, in the direction of the anti-Center region.  The disks in our simulations develop outer arcs of material generated in association with each disk crossing. These evolved outer wrappings are loosely wound and resemble rings. One of the predicted arcs, at about 10 kpc from the Sun, is reminiscent of the low-latitude Milky Way feature known as the Monoceros ring. Though the Monoceros ring is often considered to be the leftover tidal stream from a now-defunct dwarf satellite galaxy \\citep{penarrubia_etal05}, some observational evidence has suggested that the Monoceros ring could be a feature of the Milky Way itself \\citep{momany_etal06}. Previous theoretical work has suggested that a past encounter with some previously unidentified massive satellite could have produced the Monoceros ring as the outcome of a disk impact \\citep{kazantzidis_etal08}, and that these structures are long-lived during the impact itself \\citep{purcell_etal09}. We specifically identify the Sgr progenitor as the likely candidate for the impact that molded the Monoceros ring from the Milky Way disk, as the induced spiral arms detached from the outer Galactic plane and began to oscillate vertically over a range of $5-10$~kpc (as shown in the right panel of Figure~\\ref{fig:spiral}).  \\begin{figure}[t] \\begin{center} \\includegraphics[width=5.2in]{purcellc_fig3.eps} \\end{center} \\caption{{\\em Left} panel: the Toomre stability parameter Q and the critical wavelength associated with scales most susceptible to instability. Where necessary, we set the mode number $m=2$ for two-armed spiral and bar structure. Note that the initial Milky Way model, as well as both endstate disks from the {\\em Light Sgr} and {\\em Heavy Sgr} infall models, are globally stable with $Q > 1$ at all radii where disk particles dominate the stellar density. Small-scale perturbations with wavelengths less than roughly 2-3 kpc travel easily through all of these disks. {\\em Right} panel: the phase (or shape function) $\\phi$ as a function of radius, during the second pericentric approach/impact of the {\\em Heavy Sgr} satellite, and annotated regionally with values of radial wavenumber $k = d\\phi / dR$. Near the radius of disk impact (shaded area), a leading density disturbance emerges due to the gravitational perturbation of the interloping subhalo, and this feature is quickly sheared by differential rotation into the generally-trailing spiral structure of the post-infall disk as time evolves.} \\label{fig:modes} \\end{figure}  In order to quantify the presence of spiral arms in the resultant disk of the fiducial infall model, we have examined the Fourier amplitudes of the final density distribution, finding significant power in the $m=2$ and $m=4$ modes (corresponding to two-armed and four-armed morphologies respectively), with comparatively negligible power in the first two odd modes relative to axisymmetry. This is broadly consistent with analysis of satellite impacts onto the Milky Way \\citep{chakrabarti_blitz09}, which finds that satellites with mass ratio $\\sim1:50$ and orbital pericenter distances $\\sim 20$ kpc do induce power in the first few Fourier modes, with comparatively greater amplitude in the $m=2$ mode. Analysis of the Toomre stability $Q$ parameter \\citep[which similarly assesses axisymmetric disk instability, again see {\\em e.g.}][]{binney_tremaine08} and associated critical wavelength $\\lambda_{\\mathrm{crit}}$ indicates that our fiducial Milky Way model is globally stable against long-wavelength perturbations, but is susceptible to short-wavelength modes on small scales, as demonstrated by Figure~\\ref{fig:modes}. ",
        "conclusions": ""
    },
    "1402/1402.5109_arXiv.txt": {
        "abstract": "Much of the evolution of galaxies takes place in groups where feedback has the greatest impact on galaxy formation. By using an optically selected, statistically complete sample of 53 nearby groups (CLoGS), observed in Radio (GMRT) and in X-rays (Chandra \\& XMM-Newton), we aim to characterize the radio-AGN population in groups and examine their impact on the intra-group gas and member galaxies. The sensitivity to older electron populations at 240 MHz and the resolution of 610 MHz is the key to identify past and current AGN activity. Here we will present first results from three different galaxy groups analysed so far. We report an age of $\\sim$18 Myr for the radio source 3C270 in NGC~4261 implying that the expansion was supersonic over a large fraction of its lifetime. In NGC~1060 we detect a small-scale (20''/7.4 kpc) jet source with the spectral index of $\\alpha_{240}^{610}$=0.9 indicating a remnant of an old outburst. Lastly in NGC~5982 the 610 \\& 235 MHz observations find a radio point source in the central AGN with no evidence of jets or lobes and diffuse emission from the disks (star formation). ",
        "introduction": "\\label{s:intro}  Much of the evolution of galaxies takes place in groups where the balance between hot and cold gas becomes significant for the investigation of the hot intra-group medium (IGM)/active galactic nuclei (AGN) connection. The transfer of energy though between AGN and the IGM still remains an unresolved issue. Since most of the present-day galaxies and baryonic matter in the Universe reside in groups (\\citealt{Eke2004}, \\citealt{Mulchaey2000}) the study of feedback across a wide variety of nearby groups is fundamental, if we are to understand how it has influenced the thermal history of IGM. Such a study requires multi-frequency radio data and high-quality X-ray observations in order to provide insight into the physical mechanisms of the energy injection.  We define an optically selected, statistically complete sample of 53 nearby groups (Complete Local-Volume Groups Sample project), observed in Radio using GMRT in dual-frequency 610/240 MHz and in X-rays using the Chandra and XMM-Newton satellite observatories, very carefully screened to exclude systems which are uncollapsed (see \\citealt{OSullivan2014} for more details of sample selection). Using archival radio data and new observations, we target the characterization of the radio-AGN population in groups of galaxies and the examination of their impact on the intra-group gas and member galaxies.  The examination of low-frequency radio emission is particularly important in order to identify past as well as current AGN activity, combining the good spatial resolution of the GMRT at 610 MHz and the sensitivity to older electron populations at 240 MHz. Multi-frequency observations also allow the production of spectral index maps. Using a variety of systems which host different types of AGN, the possible mechanisms of AGN feedback heating in groups, the length of the AGN duty cycle and the relationship between environment and AGN jet/lobe properties can be constrained.  Here results from three different systems that exhibit the apparent diversity of the sample will be presented; We examine a group dominated by a single powerful central source where large-scale jets inflate radio lobes (NGC4261), a group combining weak AGN and star-formation (NGC5985) and a recently confirmed merger group with multiple AGN and star formation detections (NGC1060). ",
        "conclusions": ""
    },
    "1402/1402.5279_arXiv.txt": {
        "abstract": "Recent observational and theoretical investigations have emphasised the importance of filamentary networks within molecular clouds as sites of star formation. Since such environments are more complex than those of isolated cores, it is essential to understand how the observed line profiles from collapsing cores with non-spherical geometry are affected by filaments. In this study, we investigate line profile asymmetries by performing radiative transfer calculations on hydrodynamic models of three collapsing cores that are embedded in filaments. We compare the results to those that are expected for isolated cores. We model the five lowest rotational transition line (J = 1-0, 2-1, 3-2, 4-3, and 5-4) of both optically thick (HCN, HCO$^+$) as well as optically thin (N$_2$H$^+$, H$^{13}$CO$^+$) molecules using constant abundance laws. We find that less than 50\\% of simulated (1-0) transition lines show blue infall asymmetries due to obscuration by the surrounding filament. However, the fraction of collapsing cores that have a blue asymmetric emission line profile rises to 90\\% when observed in the (4-3) transition. Since the densest gas towards the collapsing core can excite higher rotational states, upper level transitions are more likely to produce blue asymmetric emission profiles. We conclude that even in irregular, embedded cores one can trace infalling gas motions with blue asymmetric line profiles of optically thick lines by observing higher transitions. The best tracer of collapse motions of our sample is the (4-3) transition of HCN, but the (3-2) and (5-4) transitions of both HCN and HCO$^+$ are also good tracers. ",
        "introduction": "\\begin{figure*} \\centering \\begin{minipage}{\\textwidth} \\centering \\includegraphics[width=\\textwidth]{01_multi_dust} \\caption{Core A, B and C in dust emission at 850 $\\mu$m and at $i$ = 0$\\degr$ and $\\phi$ = 0$\\degr$ \\citep{Smith2012a}.} \\label{fig_01_cores} \\end{minipage} \\end{figure*}  While it has long been known that stars form in dense cores of gas within molecular clouds \\citep[e.g][]{Myers1985,Cernicharo1991,Lada1993,Larson2003} it has only recently become appreciated that these cores are intrinsically connected to their environment through a network of filaments \\citep[e.g][]{Menshchikov2010,Arzoumanian2011,Peretto2012}. Our understanding of the physics and predicted observable properties of isolated cores is well developed, but the effects of surrounding filaments on observed core properties is still at an early phase of investigation. This paper is the third in a series that attempts to investigate the predicted line profiles of collapsing cores embedded within filaments.  Blue asymmetries in the line profiles of optically thick tracer molecules appear to be a good indicators of gas infall motions within isolated cores \\citep{Zhou1991,Zhou1992,Walker1994,Myers1996}. This effect arises because in optically thick species an observer sees the outer layers of the core on the red side of the line profile and the inner layers on the blue. Using this effect several observers have tried to find nearby pre-stellar cores with infall motions. \\citet{Gregersen1997,Lee1999,Gregersen2000} and \\linebreak \\citet{Sohn2007} investigate samples of Class 0 and pre-stellar cores, which are potentially collapsing, and find that less than 65\\% of the observed line profiles are blue asymmetric. In general, blue line profiles are clearly detected in only about 40\\% of the cores. It is crucial to understand how many pre-stellar cores are collapsing in order to understand how rapid and dynamic the process of star formation is. However, it is not clear if a blue asymmetries should always be expected from a core embedded within a filamentary cluster of protostars. For a more detailed description of the theory behind the blue line asymmetry and the observational literature see, e.\\ g., \\citet{Evans1999} and \\citet{Smith2012a}.  In the first paper of our series about the line profiles of cores within clusters \\citet{Smith2012a} (hereafter S12) investigate three filamentary cores, that are part of a dense giant molecular cloud (GMC) simulation, from 14 different viewing angles. Though all the model cores are collapsing, less than 50\\% of the \\linebreak HCN (J = 1-0) profiles exhibited blue asymmetries. In \\citet{Smith2013} this work is extended to look at massive star forming regions. It is found that the line profiles frequently do not have a central self-absorption dip as the collapsing regions do not always have surrounding static envelopes. Furthermore, in massive star forming regions multiple line components in optically thin species such as N$_2$H$^+$ are visible due to internal substructure. The observation of such line profiles can be used to test models of massive star formation. In this paper we seek to extend the analysis of S12 to consider additional line transitions and chemical species.  In S12 we consider the base (1-0) transitions for the sake of simplicity, since we already have a large number of variables. However, one of the reasons the success rate of the blue asymmetry is so low in the study is that the filaments in which the cores are embedded are also brightly emitting in HCN (1-0) which obscures the emission from the core. Species and transitions with a greater critical density than HCN (1-0) will only become bright at greater densities and temperature, and are consequently more likely to be visible in the core but not the surrounding filament. It would be highly advantageous to an observer surveying dense cores to know in advance which species and line transitions are most likely to trace infalling gas in pre-stellar cores embedded within dense networks of filaments.  This is the question we attempt to answer in this paper: What is the most reliable tracer of core collapse in dense clusters? In Section \\ref{sec_methods} we describe the technical tools and set-ups which have been used for this study. The results are presented and discussed in Section \\ref{sec_results}. Section \\ref{sec_conclusion} provides a summary and final conclusions. ",
        "conclusions": "The goal of this study is to identify an optimal, reliable tracer for collapsing cores in dense, filamentary clusters. For this, we continue the analysis of \\citet[S12]{Smith2012a} and investigate the question of whether we can expect blue infall line profiles in irregularly shaped, collapsing cores which are embedded in filaments.  S12 select three cores out of a Giant Molecular Cloud simulation, because they are irregular, deeply embedded in filaments, already collapsing, and about to form stars in the near future. This creates a sample of real star-forming regions with complex density and velocity patterns. However, the analysis of S12 is limited to only one optically thick and one optically thin transition. We address this limitation by modelling the (1-0) -- (5-4) transitions of two optically thick (HCN, HCO$^+$) and two optically thin (N$_2$H$^+$, H$^{13}$CO$^+$) species from twelve different line of sights. We assume constant abundances and ignore depletion effects. We use the normalised velocity difference, $\\delta v$, as defined in equation (\\ref{equ_def_deltav}) to classify the line profile asymmetries of optically thick tracers. We also introduce criteria for ambiguous spectra, see equations (\\ref{def_amb_cr1a}) and (\\ref{def_amb_cr2a}), that sort out the cases without clearly recognisable asymmetries. We study the number density, velocity and optical thickness distribution to localise the origin of line profiles features.  We find that we can expect blue asymmetric line profiles in irregular cores, as we can in spherical cores. Although less than 50\\% of simulated line profiles are blue asymmetric in the (1-0) transitions, the fraction of blue line profiles increase when going to higher transitions of the optically thick tracers. We obtain our best results in the (4-3) transitions of HCN and HCO$^+$ where about 90\\% of the line profiles are blue asymmetric. The success rates of the (3-2) and (5-4) transitions are only marginally lower, but they contain physical or observational disadvantages. The (3-2) transition is more affected by the surrounding filaments, whereas the (5-4) transition is fainter (see Table \\ref{tab03_Tmb}) and less optically thick. Thus, the (4-3) transition offers the best combination of detectability of blue line profiles and visibility above typical noise levels. This makes it the most reliable tracer for dense core gas motions in star-forming clusters.  We confirm this result by introducing classification criteria which take typical noise levels of such regions into account. They are described by the equations (\\ref{def_amb_cr1b}) and (\\ref{def_amb_cr2b}). This method reduce the maximal fraction of blue asymmetric line profiles to 60\\% -- 75\\% (depending on the used noise threshold). These results are in good agreement with observational studies of, for example, \\citet{Gregersen1997}, \\citet{Gregersen2000} and \\citet{Sohn2007}. Still, the (4-3) transitions produce the most reliable results.  We find a small but significant number of spectra that show subtle shoulder-like features in the higher (4-3) and (5-4) transitions as opposed to a clear double-peaked line profile. In some cases such features may require prohibitively high signal-to-noise to be detected in real observing conditions. We suggest that the brighter (3-2) transitions may be a better avenue in such conditions. However, we urge caution in interpreting the (3-2) line profile, as discussed above.  Furthermore, we investigate the origin of `false' line profile asymmetries by using the ODS. These are localised in the filaments around the cores where the gas is dense enough to become optically thick and hide the core behind. Since higher transitions arise from denser material they become more optically thin within the filaments. Thus, the cores and their blue line profiles become detectable. In addition, the ODS show that the optically thin tracers may become optically thick in the lower transitions. Since they are needed as reference lines in the $\\delta v$ analysis we also use higher transitions of the optically thin tracers to improve the analysis. The (3-2) transitions turn out to be the best choice, since they are usually optically thin, bright and have central frequencies close to the (4-3) transitions of the optically thick tracers.  We briefly explore the effect of depletion upon the line profiles and find that depletion in dense gas may change the details of individual spectra. However, the overall percentage of blue, red and ambiguous line profiles is largely unchanged. In this study we focus on the line profiles of collapsing clustered cores resulting purely from the complex gas dynamics and geometries within these regions. This allows us to isolate the effects of geometry upon line profiles in a relatively simple scenario that can be more easily understood. However, in future work it would be beneficial to combine the effects of geometry and dynamics with a more detailed model of chemical inhomogeneities within the dense cores. \\smallskip  In summary, we can expect to detect blue infall line profiles in irregular, collapsing cores. Dense filaments around the cores along the line of sight cause non-blue asymmetric line profiles of optically thick tracers and non-Gaussian line profiles of optically thin tracers. Using higher transitions can avoid these effects. We obtain the best results with the (4-3) transitions of the optically thick tracers while using the (3-2) transition of optically thin tracers as reference lines. Future observations with, e.\\ g., ALMA will be able to resolve the shapes of line profile, including self-absorption signatures and shoulder-like features, more precisely."
    },
    "1402/1402.2490_arXiv.txt": {
        "abstract": "We study the impact of ionospheric disturbances on the Earth's environment caused by the solar events that occurred from 20 April to 31 May 2010, using observations from the {\\it Mexican Array Radio Telescope} (MEXART). During this period of time, several astronomical sources presented fluctuations in their radio signals. Wavelet analysis, together with complementary information such as the vertical total electron content ({\\it vTEC}) and the {\\it Dst} index, were used to identify and understand when the interplanetary scintillation (IPS) could be contaminated by ionospheric disturbances (IOND). We find that radio signal perturbations were sometimes associated with IOND and/or IPS fluctuations; however, in some cases, it was not possible to clearly identify their origin. Our Fourier and wavelet analyses showed that these fluctuations had frequencies in the range $\\approx$\\,0.01\\,Hz -- $\\approx$\\,1.0\\,Hz (periodicities of 100\\,s to 1\\,s). ",
        "introduction": "\\label{intro}  Large-scale solar-wind disturbances in the interplanetary medium (IPM) can, and frequently do, distort radio-wave fronts coming from compact radio sources producing the so-called interplanetary scintillation (IPS) \\citep{hew64,hew86}. It is known that radio sources with an angular diameter smaller than 1 arcsec present  interplanetary scintillation at frequencies $>$\\,0.1\\,Hz  \\citep{hew87}; though, depending on the observing wavelength, IPS can be present beyond 10\\,Hz \\citep{mil76,man90,man94,man10}. IPS on a large number of compact radio sources has been observed using the {\\it Ooty Radio Telescope} (ORT), {\\it Solar-Terrestrial Environment Laboratory} (STEL) and {\\it Mexican Array Radio Telescope} (MEXART), providing insight on solar-wind properties \\citep{man90,man94,jul10}.  Solar events can affect the plasma of the Earth's ionosphere either directly, through X-ray and/or EUV radiation from flares as it reaches Earth, or indirectly, through the perturbation of the magnetospheric electric field upon the arrival of an interplanetary CME \\citep{Tsu09}. The Earth's ionosphere is considered a dispersive medium for radio waves. Its refractive index is a function of the radio-wave frequency, the electron density, and the intensity of the Earth's magnetic field, {\\bf B}. Its perturbation can be an important source of error for the signals of the {\\it Navstar Global Positioning System} (GPS) satellites and other positioning systems like {\\it Galileo} (European Union), {\\it Glonass} (Russia), and {\\it Compass} (China). The error in the signals is proportional to the integrated electron density along the signal path, TEC, and inversely proportional to the square of the carrier phase frequency, $\\tau \\propto TEC/\\nu ^2$ \\citep{lan96,kom97}. The signals from GPS satellites must travel through the Earth's ionosphere in their way to the GPS receivers (on the Earth's surface) in the L-band: $L_1=1575.42$\\,MHz and $L_2=1227.60$\\,MHz. Although there is another L-band, $L_3=1381.05$\\,MHz, GPS users cannot use it \\citep{eri01}.  Several works have shown the importance of the ionospheric interference on the IPS \\citep{gapp82,pur87,tap84,woa95,luc96,jac98,per08,shi10}. Part of this contamination, or ionospheric scintillation (IONS), can represent a serious problem in IPS studies. In fact, the presence of IONS can lead to a misinterpretation of the IPS; however, there are techniques to remove the ionospheric scintillation part. \\citet{car11} have explored another possible data contamination by ionospheric effects. They have shown an example related to a solar event on 15 December 2006, where the fluctuations observed in the signal of the radio sources exhibited strong contamination by ionospheric disturbances (IOND).  The behavior shown by the radio signal was consistent with the fluctuations observed from the signal emitted by polar satellites. The authors suggested that not only the total electron content (TEC) and the geomagnetic equatorial index, {\\it Dst}, could be used to characterize the origin of such fluctuations, but also the possibility of removing the contaminating frequencies from the data to study separately the IPS or the IOND.  Since we have observed fluctuations in the radio signal of sources transiting the MEXART observatory, the purpose of our study is to understand the nature of these fluctuations, as well as answering the following questions: Are the fluctuations intrinsic to quiet solar-wind conditions? When are these fluctuations related to ionospheric disturbances? How do these fluctuations quantitatively affect the IPS? To answer these questions, we have analyzed the radio signal of sources observed with the MEXART. We have explored all observations to find if the ionospheric contribution is always present, even in the absence of solar events. In Section \\ref{obs}, we present the information related to the MEXART observations for the period between April and May 2010. Section \\ref{dat_an} shows the data analysis applied and the results using wavelet analysis applied to the MEXART data. Finally, in Sections \\ref{dis} and \\ref{con} we present the discussion and conclusions of our work, respectively. ",
        "conclusions": "\\label{con}  The conclusions of our work can be summarized as follows:  \\begin{itemize} \\item Our study incorporates the use of wavelets together with complementary information provided by the {\\it vTEC} and {\\it Dst} index. This is the first time these tools have been applied to MEXART data and allow a better understanding of cases when  the IPS may be contaminated by IOND. \\item Within the period of April--May 2010, several radio sources presented fluctuations in their radio signal registered with the MEXART. Using wavelets we found the periods/frequencies that characterize these fluctuations. We found that the perturbations in the signal were sometimes associated with IOND and/or IPS (see for instance Figure \\ref{fig4} top left, Figure \\ref{fig5} top left and Figure \\ref{fig8} top right). In this context the {\\it vTEC} can help to establish the cases in which the IPS is actually contaminated by IOND. Furthermore, in the example shown in Figure \\ref{fig8} (top left), the wavelet analysis indicates that the ionospheric scintillation is associated with quiet conditions, and in Figure \\ref{fig5} (top left) and Figure \\ref{fig6} (top right) the wavelet analysis shows ionospheric disturbances associated with the geomagnetic storm. \\item In addition, Fourier spectra show that these fluctuations have frequencies between $\\sim$\\,0.01 Hz to $\\sim$\\,1.0 Hz (periods of 100\\,sec to 1\\,sec, respectively) for both IOND and IPS. We found that there is a tendency for the frequencies close to $\\sim$\\,1.0 Hz to be more likely associated with IPS, whereas lower frequencies tend to be related to IOND. However, in the cases where both IPS and IOND occur, we could not distinguish between them. \\item In the examples that neither the Dst index nor the {\\it vTEC} showed any significant change, we can say that the fluctuations are not related to global ionospheric effects but could be probably related to IPS. We observed that when the {\\it vTEC} is perturbed and  solar-origin effects can be assumed, then the radio signal can be contaminated by ionospheric perturbations. This is probably the case observed by \\citet{car11}, where the fluctuations could be attributed to IOND and/or IPS. \\item The solar-minimum period has allowed us to characterize better the radio fluctuations of astronomical sources observed with the MEXART, during this period the ionosphere can contribute strongly. This is especially true for those sources that having large elongations (around 90$^{\\circ}$, for instance) present fluctuations in their signal. However, more work is needed to characterize the intrinsic fluctuations in the radio signal. \\end{itemize}"
    },
    "1402/1402.1493_arXiv.txt": {
        "abstract": "We investigate baryon effects on the halo mass function (HMF), with emphasis on the role played by AGN feedback. Halos are identified with both Friends-of-Friends (FoF) and Spherical Overdensity (SO) algorithms. We embed the standard SO algorithm into a memory-controlled frame program and present the {\\bf P}ython spher{\\bf I}c{\\bf A}l {\\bf O}verdensity code --- {\\small PIAO} (Chinese character: \u6f02).  For both FoF and SO halos, the effect of AGN feedback is that of suppressing the HMFs to a level even below that of Dark Matter simulations. The ratio between the HMFs in the AGN and in the DM simulations is $\\sim 0.8$ at overdensity $\\Delta_c=500$, a difference that increases at higher overdensity $\\Delta_c=2500$, with no significant redshift and mass dependence. A decrease of the halo masses ratio with respect to the DM case induces the decrease of the HMF in the AGN simulation. The shallower inner density profiles of halos in the AGN simulation witnesses that mass reduction is induced by the sudden displacement of gas induced by thermal AGN feedback. We provide fitting functions to describe halo mass variations at different overdensities, which can recover the HMFs with a residual random scatter $\\simlt 5$ per cent for halo masses larger than $10^{13} \\hMsun$. ",
        "introduction": "\\label{i}  The halo mass function (HMF hereafter) is a unique prediction of cosmological models of structure formation. The evolution of the HMF traced by galaxy clusters has been recognized since a long time as a powerful tool to trace the growth of cosmic structures and, therefore, to constrain cosmological parameters \\cite[see][for reviews, and references therein]{Rosati2002,Allen2011}. In particular, cosmological applications of the HMF require to know its shape and evolution to a high precision, in order to fully exploit its potential as a cosmological tool to be applied to ongoing and future large surveys of galaxy clusters \\citep[e.g.][]{Wu2010,Murray2013}.  N--body simulations covering wide dynamic ranges are nowadays providing rather accurate calibration of the mass function of Dark Matter (DM) halos \\citep[e.g.][]{Jenkins2001, Reed2003, Reed2007, Reed2013,Lukic2007,Tinker2008,Crocce2010,Courtin2011, Bhattacharya2011,Angulo2012,Watson2013}. Various extensions of the standard $\\Lambda$CDM cosmology model, such as coupled dark energy models \\citep[e.g.][]{Cui2012b, Baldi2012}, modified gravity models \\citep[e.g.][]{Schmidt2009,Zhang2013,Puchwein2013}, non-Gaussian initial conditions \\citep[e.g.][]{Grossi2009,Pillepich2010}, massive neutrinos \\citep[e.g.][]{Brandbyge2010,Ichiki2012,Costanzi2013}, Warm Dark Matter \\citep[e.g.][]{Schneider2013,Angulo2013}, have been studied using numerical simulations, and their effect on the HMF has been investigated.  A crucial aspect in the calibration of the HMF is related to the algorithm used to identify halos, the two most widely used being the Friend of Friend (FoF) one and the Spherical Overdensity (SO) one. The choice of a specific algorithm clearly impacts both the number of identified halos and their mass \\citep[e.g.][; see also \\citealt{Knebe2011} for a detailed comparison between different halo finders]{White2001, White2002, Lukic2009, More2011, Watson2013}.  All the above mentioned HMF calibrations are based on N--body simulations that follow the evolution of a collisionless DM fluid. On the other hand, the presence of baryons is known to add subtle but sizeable effects on halo formation and internal structure, whose details also depend on the physical processes included in the numerical treatment of the baryonic component, such gas cooling, star formation and energy feedback from supernovae (SN) and Active Galactic Nuclei (AGN) \\citep[e.g.][and references therein]{KravtsovBorgani2012}.  A number of studies based on cosmological hydrodynamical simulations have been recently carried out to analyse in detail the effect of baryonic processes on different properties of the total mass distribution, such as the power spectrum of matter density fluctuations \\citep[e.g.]{Rudd2008, Daalen2011, Casarini2012}, the halo correlation functions \\citep[e.g.][]{Zhu2012, Daalen2013}, the halo density profiles \\citep[e.g.][]{Duffy2010, Lin2006} and concentration \\citep[e.g.][]{Rasia2013, Bhattacharya2013}, and the HMF \\citep[e.g.][]{Stanek2009, Cui2012a, Sawala2013, Martizzi2013, Cusworth2013, Balaguera2013, Wu2013}.  As for the effect of non--radiative hydrodynamics, the presence of baryons has been shown to induce a slight increase of the HMF \\citep[][hereafter Paper I]{Cui2012a}. When extra--heating is included, \\cite{Stanek2009} found instead a decrease in the HMF. As for the effect of radiative cooling, star formation and SN feedback, different groups consistently found an increase of the HMF, an effect that is more evident in the high--mass end \\citep{Stanek2009,Cui2012a,Martizzi2013}. On the other hand, \\citet{Sawala2013} found that efficient SN feedback produces an opposite effect in low--mass halos.  On the other hand, a number of analyses have shown in the last years that including AGN feedback in cosmological simulations provides populations of galaxy clusters in better keeping with observational results \\citep[e.g.][]{Puchwein2008, ShortThomas2010, Fabjan2010, McCarthy2011, Planelles2013, LeBrun2014}. When the AGN feedback is included, different results were found by \\citet{Martizzi2013} and \\citet{Cusworth2013}. The former showed that the HMF with AGN feedback is higher than the fitting function from \\citet{Tinker2008}, while the latter predicted a lower HMF compared to the same fitting function. However, their implementation of the AGN feedback differ.  \\cite{Martizzi2013} described AGN feedback by computing explicitly gas accretion rates onto super-massive black holes (SMBHs) included as sink particles in simulations that also include radiative cooling, star formation and SN feedback \\citep[e.g.][]{Springel2005,BoothSchaye2009}.  \\cite{Cusworth2013} included AGN feedback by computing the associated feedback energy from the semi--analytic model of galaxy formation by \\cite{Guo2011}, without including radiative cooling and assuming zero mass for gas particles, so that no back-reaction of baryons on the DM distribution is allowed.  In this paper we extend our previous analysis of baryonic effects on the HMF, presented by \\citetalias{Cui2012a}, by also including in our simulations the effect of AGN feedback. We directly compare the HMF obtained from DM--only simulations to those produced by radiative hydrodynamical simulations both with and without AGN feedback, using exactly the same initial conditions, mass and force resolutions. The plan of the paper is as follows. In section \\ref{simulation}, we present the simulations analysed in this paper. Section \\ref{halo} is devoted to the description of the halo identification methods. In section \\ref{results} we present the results of our analysis and describe in detail the differences in the HMF induced by different feedback models. Our results are discussed and summarised in Section \\ref{concl}. ",
        "conclusions": "\\label{concl}  Based on a set of large--scale N-body and hydrodynamical simulations, we presented an analysis of the baryon effects on the halo mass function, with emphasis on the role played by gas accretion onto super-massive black holes (SMBHs) and the ensuing AGN feedback. As such, this analysis extends our previous one, presented in \\citetalias{Cui2012a}, to the case in which one accounts for a suitable model for AGN feedback that regulates star formation within massive halos, thereby improving the degree of realism of the simulated galaxy clusters and groups. We compared three simulations: a first one including only collisionless Dark Matter; a second one including hydrodynamics with radiative physics and supernova (SN) feedback; a third one which also includes AGN feedback. Based on these simulations, we analyse how the halo mass distribution reacts to overcooling and, in the presence of AGN feedback, to the sudden displacements and expulsion of large amount of gas.  We use both Friends-of-Friends (FoF) and Spherical Overdensity (SO) algorithms to identify halos. FoF halos are identified using a standard (on-the-fly) FoF finder with a linking length $b = 0.16$. As for SO halos, they are identified using an efficient memory-controlled python code --- {\\small PIAO}, in which halos are not allowed to overlap. We focus on three overdensities $\\Delta_c = 2500, 500, 200$ for the SO halos.  The main results of our analysis are summarised as follows. \\begin{description}  \\item[1.] Including AGN feedback in hydrodynamical simulations causes a reduction of the halo mass function (HMF) with respect to the DM--only case, by an amount which depends on the overdensity within which halo mass is measured. This effect amounts to about 20 per cent for overdensity $\\Delta_c=500$, with no evidence of mass and redshift dependence, and increases at higher overdensity. Therefore, our model of AGN feedback reverses the effect of radiative physics with no efficient feedback, which instead leads to an increase of the HMF. \\item[2.] The baryonic effects on the HMF can be traced to the effect that different feedback models have on the halo density profiles and total masses. In the absence of AGN feedback, we confirm that halo density profiles steepen as a consequence of adiabatic contraction associated to overcooling, thus causing an increase of halo masses. AGN feedback generates shallower density profiles and a corresponding decrease of halo masses. This effect is caused by the improved regulation of overcooling associated to AGN feedback and to the sudden displacement of large amount of gas, which makes the gravitational potential responding with an expansion. The relative decrease of halo masses is larger in smaller objects, where AGN feedback is relatively more efficient, and at higher overdensities; for $\\Delta_c=500$, it amount to $\\sim 20$ per cent at $\\log M_{500}=12.5$, while becoming negligible for the largest halos found in our simulation box, with $\\log M_{500}\\simeq 15$. Interestingly, this effect is independent of redshift, at least up to $z=1$ where we have a large enough statistics of massive halos. \\item[3.] We provide a mass--dependent fit to the halo mass variations induced by baryonic effects. After applying this model for the mass correction to the HMF obtained from the DM--only simulation, we recover the HMF from hydrodynamical simulations, up to a random scatter of $\\simlt 5$ per cent for halos more massive than $10^{13} \\hMsun$, with no significant bias and independently of redshift. \\end{description}  Our analysis demonstrates that baryon effects can cause sizeable variations of the HMF and, therefore, affects the measurement of cosmological parameters from redshift number counts of galaxy clusters. For instance, recent results from the Planck Collaboration \\citep[cf. also \\citealt{Spergel2013}]{PlanckXXI} highlights that the cosmological model preferred by CMB analysis over-predicts the number of clusters expected in the SZ Planck Cluster Survey by about 50 per cent. Interestingly, this tension would be relaxed, although by a small amount, if the HMF calibrated with our AGN simulation is used to predict SZ cluster number counts.  In a recent paper, \\cite{Khandai} also investigated the effect of including AGN feedback on the HMF. Since they considered simulation boxes smaller than our, with size of 100$\\hMsun$, they better probed the low-mass end of the HMF, while having a worse sampling of the high mass end. As a result of their analysis, they found that the FoF halo mass function can be described with a universal form to a reasonable accuracy, even after accounting for baryon effects.\\footnote{After the submission of our paper, a paper by \\cite{Velliscig14} appeared on the arXiv, which also included an analysis of baryon effects on halo masses and HMF, when AGN feedback is also included. Using box size and resolution quite similar to those of our simulations, they confirmed our results on the effect of AGN feedback on the HMF.}  Numerical convergence in the calibration of the HMF from purely collisionless simulations could in principle be reached by ``brute force'' (i.e. increasing the dynamic range accessible and the gravitational force integration accuracy). The same may not be true when the effects of baryons are included. In fact, our analysis shows that baryons can generate variations of the HMF with respect to the DM case which goes in different directions, depending on the nature and, possibly, on the numerical implementation of feedback. In this respect, confidence in the calibration of baryon effects in the HMF is strictly related to the level of agreement between observational results and model predictions for a variety of properties of galaxy clusters. Current implementations of AGN feedback in cosmological simulations produce cluster populations with an increased degree of realism \\citep[e.g.][ and references therein]{Puchwein2008,Dubois2012, Planelles2013, Planelles2014, LeBrun2014}. As an example, we compare in Appendix \\ref{A:sbf} two basic properties involving baryons in clusters, namely the stellar mass fraction and the total baryon mass fraction, with observational results. The good agreement between the AGN simulation and observational results in quite encouraging. Still, it is fair to admit that important tensions still exist between a number of observed and simulated cluster properties, such as the thermal structure of cool cores and the properties of the Brightest Cluster Galaxies \\citep[e.g.][ cf.  \\cite{Kravtsov2014}]{Dubois2011, Ragone2013, Planelles2014}. There is no doubt that providing an HMF, that accounts for the inclusion of the baryon physics at the percent level of accuracy required by the next generation of large--scale cluster surveys, still requires substantial work and, ultimately, a precise numerical description of galaxy formation in a cosmological framework."
    },
    "1402/1402.6566_arXiv.txt": {
        "abstract": "There is a $\\sim 150\\kms$ discrepancy between the measured  motion of the Local Group of galaxies (LG) with respect to the CMB and the linear theory prediction  based on the gravitational force field of the large scale structure in full-sky redshift surveys. We perform a variety of tests which show that  the LG  motion  cannot be recovered to better than $150-200\\kms$ in amplitude and within a $\\approx10^\\circ$ in direction. The tests rely on catalogs of mock galaxies identified in the Millennium simulation using semi-analytic galaxy formation models. We compare these results to the $K_s=11.75$ Two-Mass Galaxy Redshift Survey, which provides the deepest, widest and most complete spatial distribution of galaxies available so far. In our analysis we use a new, concise relation for deriving the LG motion  and bulk flow from the true distribution of galaxies in redshift space.  Our results show that the  main source of uncertainty is the small effective depth of  surveys like the 2MRS that prevents  a proper sampling of the large scale structure beyond $\\sim100\\hmpc$. Deeper redshift surveys are needed to reach the  ``convergence scale\" of  $\\approx 250\\hmpc$ in a $\\Lambda$CDM universe. Deeper survey would also mitigate the impact of the ``Kaiser  rocket\" which, in a survey like 2MRS, remains a significant source of uncertainty. Thanks to the quiet and  moderate density environment of the LG, purely dynamical uncertainties of the linear predictions are subdominant at the level of $\\sim 90\\kms$. Finally, we show that deviations from linear galaxy biasing and shot noise errors provide a minor contribution to the total error budget. ",
        "introduction": "\\label{sec:int}  The group of galaxies containing M31, the Milky Way (MW) and about a dozen other, much smaller,  galaxies (excluding satellites) within $\\sim 1.4$ Mpc form a bound system which is detached from the general cosmic expansion \\citep{yst77}. This Local Group (LG) of galaxies resides in mildly over-dense region characterized by a remarkably small velocity shear. Just like any other cosmological object, the LG  is expected to move with a non-vanishing velocity relative to the general expanding background. The best approximation to the frame of reference defined by the cosmological background is undoubtedly based on temperature maps of the Cosmic Microwave Background (CMB). The high degree of dipole anisotropy in the temperature of the CMB on the sky \\citep{1993ApJ...419....1K, 1996ApJ...473..576F, 2009ApJS..180..225H} is interpreted as a Doppler boosting resulting from our motion through a  highly isotropic thermal CMB photons. This interpretation  has recently been reinforced  by the detection of the corresponding   modulation and aberration of the CMB fluctuations observed by the Planck satellite \\citep{2013arXiv1303.5087P}.  The observed dipole yields a very precise measurement of the Solar barycenter velocity relative to  the ``CMB frame\" of reference in which an observer at rest would measure a vanishing  dipole anisotropy. Augmented with astronomical estimation of the LG motion relative to the Sun (c.f. \\S\\ref{sec:refs}), the CMB dipole provides  $V_{\\rm lg}=627\\pm 22\\kms$ toward $(l,b)=(276^\\circ\\pm 3^\\circ,30^\\circ\\pm 3^\\circ)$ as the reference value for the LG motion relative to the CMB frame \\citep{1993ApJ...419....1K}. In the standard cosmological paradigm \\citep{Peeb80}, the LG is accelerated by the cumulative gravitational pull of the surrounding large scale structure. In linear theory, the peculiar velocity is proportional to the peculiar gravitational force field times the Hubble time, with proportionality factors depending on the background  mass density. Therefore, it is natural to ask whether the observed large scale structure, as traced by the galaxy distribution, could indeed account for the LG motion.   This issue was recently studied by \\cite{DN10} who found good agreement between the local velocity and gravitational fields, in contrast to the earlier data with an inferior velocity field, which gave irreconcilable differences between the two (\\cite{DNW96}). The earlier studies of  \\citep{1980ApJ...242..448Y,DH82} addressed the gravity versus velocity fields, with somewhat contradictory results, using either the angular positions and fluxes from photometric galaxy catalogs \\citep{yahil86,MD86,harmon87,Vill87,P88,bilicki11} or the full 3D distribution of different types of extragalactic objects ranging from the infra-red selected galaxies of the IRAS 1.2 Jy catalog \\citep{strauss92,web97,zaroubi99} and extension to fainter fluxes, the PSC$z$ catalog \\citep{schmoldt99,schmoldt99b,rr00,bp06}, optically selected galaxies \\citep{Lahav87,Lybell89,huds93}, mixed catalogs of infra-red and optical galaxies \\citep{Lahav88,dmellow04}, galaxy clusters selected from optical plates \\citep{PV91,scaramella91,scaramella94,bps96,dale99} and, finally, X-ray selected galaxy clusters \\citep{pk98,ke06}.   In this work we focus on determining how well the LG motion can be recovered from  the observed galaxy distribution via linear instability theory. Detection of significant departures from theoretical expectations require a complete understanding and characterization of all possible error sources. In the case of cosmological dipoles these are \\citep{schmoldt99} \\begin{itemize} \\item Cosmic variance from finite volume sampling. All sky galaxy surveys become significantly more dilute at larger distances, limiting the depth within which density fluctuations can be reliably probed for the reconstruction of the LG motion. \\item Shot noise from sparse sampling of the mass tracers. \\item Highly non-linear dynamical effects that spoil the relation the tight relation between the gravity and the LG motion. They include nonlinear motions and growth of density fluctuations. \\item Deviations from strict linear biasing between galaxies and the underlying mass density fluctuations. \\item Observational uncertainties and biases arising from incompleteness and selection effects in the parent objects' catalogue.  \\end{itemize}  The most appropriate  route to estimate the impact of these uncertainties and  assess the adequacy of linear theory has   been pioneered by \\cite{1991ApJ...372..394D}. It relies on  the extensive use of realistic mock galaxy catalogs extracted from N-body simulation, since they simultaneously account for non nonlinear effects as well as selection effects specific to the specific dataset. In particular, the mocks should have LG candidates residing in a  mild density region  and small velocity shear, as in the real data. Although,  one can attempt to account for  non-linear dynamical effects within a likelihood formalism by quantifying the \"decoherence\" between the gravity and the peculiar velocity fields \\citep{chodor02,chodor04,chodor08}, we opt here to rely on mock catalogs extracted from fully nonlinear $N$-body simulation. While we aim at a general discussion of the problem, we shall consider here the case of the  Two-Mass  Redshift Survey  (2MRS) \\citep{huchra12},  i.e. the the deepest nearly-all sky survey of angular positions and spectroscopic galaxy redshifts limited to $K_s=11.75$ and arguably the best sample of objects to estimate the LG motion.  The outline of the paper is as follows. A brief description of how the  LG motion has been measured in the literature is given in \\S\\ref{sec:refs}. In \\S\\ref{sec:equations}, we review the linear theory predictions and offer a new useful relation for deriving the LG motion from a given distribution of mass tracers in redshift space. In \\S\\ref{sec:mocks}  we first test  linear theory  as measured by the full dark matter distribution  and the full volume-limited galaxy distribution. We characterize the biasing relation of the  galaxy catalogs, address the reliability of linear reconstruction of the LG motion,  and assess the impact shot noise errors. In \\S\\ref{sec:dmrec}  we consider the analysis when applied to  mock 2RMS  catalogs. The results are presented in  \\S\\ref{sec:results} in which we compare the LG motions obtained from the distribution of mock galaxies to that measured directly in the N-body simulation. The impact of the so-called Kaiser effect is outlined in a separate \\S\\ref{sec:kaiser} and, finally, we end with a general discussion in \\S\\ref{sec:disc}. In the Appendix the reader will find a detailed derivations of the linear theory relations used here and farther discuss the dependence on the so-called distortion parameter $\\beta$ of peculiar velocities reconstructed from redshift space data. ",
        "conclusions": "\\label{sec:disc}   In this paper we have investigated the various error sources  in the determination of the LG motion. \\begin{itemize}  \\item It is a mistake to do the analysis in the CMB frame because this frame is only gradually reached. The LG frame gives a better estimate of the distance to the nearby mass fluctuations, as seen in the substantial deviations in the residual vector differences for $R_{out} < 70\\hmpc$  (e.g.  figure (\\ref{fig:vdm})).   \\item By far the main source of error is related to the limited depth of galaxy surveys, i.e. cosmic variance. The error in the LG motion estimated from the full dark matter density field in real space and within a radius of  $250\\hmpc$ is  $\\sim 90\\kms$. However, the dilute sampling and flux-limited nature of most avalable datasets makes them significantly shallower than this. For an all sky survey like the 2MRS, the contribution to the LG motion can be assessed reliably only within $\\sim 100\\hmpc$. At this  depth, the error in the predicted LG motion is $\\sim 200\\kms$.  \\item The velocity residuals measured in the perpendicular and parallel directions allow us to estimate the {\\it rms} angle of the LG velocity from the full sky data. We find that errors of $\\approx 10^\\circ$ are inevitable, and are not sensitive to whether the analysis is done in  redshift space or real space.   \\item Another source of uncertainty are the  errors in the dynamical reconstruction. To assess their contribution to the total error budget  we have estimated the LG motion from the full dark matter out to the largest possible outer  radius, i.e. $250\\hmpc$. Then, we have filtered out the contribution from large scale structure beyond this radius by measuring the LG motion with respect to the bulk flow of the sphere. Hence the resultant errors are entirely due to dynamical inaccuracies of linear theory. The corresponding $1\\sigma$ error, $\\sim 90\\kms$, is substantially  smaller that the typical error  in linear reconstruction of the peculiar velocity of a generic observer in the Universe. The reason for this is the strict criteria we have applied in  selecting the \"LG observer\" in the mock catalogs, aimed at  matching   the quietness and moderate density  environment of the observed LG.   Removing these selection criteria boosts the  error to $\\gtsim 300 \\kms$, consistent with previous studies \\citep{NB00}.  Nonlinear dynamical reconstruction methods \\citep[e.g.][]{sh95,cc97,f02,NB00,BEN02} can potentially reduce the dynamical error. However, because the particular environment of the LG,  errors due to linear reconstruction are subdominant compared to the total error budget discussed above.  \\item The shot noise originating from the sampling of the mass density field by a finite number of tracers, forcing the use of the selection function of the sample $\\varphi$ to deal with the magnitude limit, is another error source. In the case of the 2MRS galaxies brighter than K$_{\\rm s}=11.75$,  the error amplitude is comparable to that of the dynamical error. The {\\it rms} of the combined  dynamical and shot noise scatter is close to the error  estimated directly from the scatter of the LG motions reconstructed from our mock galaxy catalogs, leaving little room for any substantial, additional source of error.     \\item Indeed, we have identified the remaining error source with galaxy biasing. Galaxies in our mocks follow a  biasing relation that is close to linear in regions with positive densities but it significantly deviates from linearity in voids where galaxies are less abundant than expected from an extrapolation of the linear bias. Moreover, the biasing relation is non-deterministic. Its scatter is driven by Poisson noise everywhere but in high density regions, where the intrinsic scatter in the biasing relation is dominant. Nonetheless, uncertainties due to deviations from the assumptions of linear biasing are of relative insignificance as illustrated by the comparison of  the real space reconstructions from 2MRS-like mocks generated from dark matter particles and the mock galaxy catalogs.  \\item Reconstructing velocities from redshift space data is fundamentally a more challenging problem than in real space. Effects like multi-valued zones (tracers with distinct distances along the same line of sight, but with very similar redshifts) and fingers of god (spreading of galaxies in virialized regions along the line of sight) affect the reconstruction on scales larger than the traditional non-linear scale.  In order to mitigate the effect of the fingers of god, the redshifts of the mock galaxies have been computed with peculiar velocities smoothed on a scale $1\\hmpc$ scale. The $1\\sigma$ errors in the LG motion reconstructed in redshift space reconstruction is  $\\sim 20\\%$ larger than in real space.  \\item We have also assessed the \"Kaiser rocket\" effect and demonstrated that it can be  partially corrected if the selection function is well constrained by observations. Correction is easy in the case of  mock catalogs but more challenging in real datasets where the effect of galaxy evolution cannot be ignored, even within the local volume encompassed by the 2MRS  \\citep{2012MNRAS.424..472B}. A distinct signature of the \"Kaiser rocket\" effect is the overshooting of the reconstructed LG motion at large radii, ($\\gtsim 100\\hmpc$). We are not aware of any reconstruction of galaxy dipoles from real datasets taking into account this spurious growth.  \\end{itemize}   Despite the general consensus on its origin, the identification of the actual gravitational sources  responsible for the LG motion is still a matter of debate. Gravity is a long range force and it may prove futile to try to identity specific sources for the LG motion. A more useful description is  in terms of ``dipole convergence scale\", i.e. the physical distance which encompasses the matter fluctuations responsible  for generating  most of the LG motion. In the standard $\\Lambda$CDM model, the dominant contribution to the LG motion is expected to arise from mass fluctuations with a distance up to $\\sim 200-300\\hmpc$ from the LG \\citep{bilicki11}. We have seen that convergence in the mocks is achieved at a  depth of $\\approx 250\\hmpc$, consistent with the theoretical expectation. The convergence is gradual, indicating that the cumulative effect of the large scale mass density field must be invoked to account for the LG motion,  The  $150-200\\kms$ accuracy  should be regarded as a conservative estimate for the expected error in the estimates of the LG velocity. Indeed,  comparisons between the observed LG motion and the prediction from various redshift surveys have yielded consistency up to the level we see here in the mocks. We, therefore, conclude that the standard picture for  the formation of large scale structure is fully consistent with current observational data of the large scale motions."
    },
    "1402/1402.3033_arXiv.txt": {
        "abstract": "Herschel observation has recently revealed that interstellar molecular clouds consist of many filaments. Polarization observations in optical and infrared wavelengths indicate that the magnetic field often runs perpendicular to the filament. In this paper, the magnetohydrostatic configuration of isothermal gas is studied, in which the thermal pressure and the Lorentz force are balanced against the self-gravity and the magnetic field is globally perpendicular to the axis of the filament. The model is controlled by three parameters:  center-to-surface density ratio ($\\rho_c/\\rho_s$), plasma $\\beta$ of surrounding interstellar gas ($\\beta_0$) and the radius of the hypothetical parent cloud normalized by the scale-height ($R'_{0}$), although there remains a freedom how the mass is distributed against the magnetic flux (mass loading). In the case that $R'_0$ is small enough, the magnetic field plays a role in confining the gas. However, the magnetic field generally has an effect in supporting the cloud. There is a maximum line-mass (mass per unit length) above which the cloud is not supported against the gravity. Compared with the maximum line-mass of non-magnetized cloud ($2c_s^2/G$, where $c_s$ and $G$ represent respectively the isothermal sound speed and the gravitational constant), that of the magnetized filament is larger than the non-magnetized one. The maximum line-mass is numerically obtained as \\[ \\lambda_{\\rm max} \\simeq 0.24 \\Phi_{\\rm cl}/G^{1/2} + 1.66 c_s^2/G, \\] where $\\Phi_{\\rm cl}$ represents one half of the magnetic flux threading the filament per unit length. The maximum mass of the filamentary cloud is shown to be significantly affected by the magnetic field when the magnetic flux per unit length exceeds $\\Phi_{\\rm cl} \\gtrsim 3\\,{\\rm pc\\,\\mu G}\\,(c_s/190\\,{\\rm m\\,s^{-1}})^2$. ",
        "introduction": "Recently, filamentary structure in molecular clouds has attracted a great deal of attention in the context of star formation. Thanks to the high sensitivity of Herschel satellite \\citep{pilbratt2010} in infrared (IR) and sub-mm ranges, \\revII{Herschel} has found many filaments in the thermal dust emissions from molecular clouds, which include clouds inactive in star formation such as Polaris \\citep{menshchikov2010,miville2010}, as well as active ones such as the Aquila cloud \\citep{menshchikov2010}, IC 5146 \\citep{arzoumanian2011}, Vela C \\citep{hill2011}, and Rosette cloud \\citep{schneider2012}. This indicates that the molecular clouds consist of the gas filaments and the star formation process should be studied in this context.  Polarization observation of background stars beyond a molecular cloud gives the geometry of interstellar magnetic field inside and around the cloud. This is based on the fact that the dust is aligned in the magnetic field and the light obscured by the intervening aligned interstellar dusts residing in the cloud shows a polarization such as the {\\bf E}-vector of the polarization is parallel to the interstellar magnetic field. From the near IR (J, H, and Ks bands) imaging polarimetry of Serpens South Cloud, \\citet{sugitani2011} have found a well-ordered global magnetic field perpendicular to the main filament. They also found that small-scale filaments seem to run along the magnetic field. Even in the Taurus dark cloud, optical and near IR polarimetry \\citep{moneti1984} indicates that the global magnetic field seems perpendicular to the major axis of the clouds. B211 and B213 filamentary clouds run in a direction perpendicular to the magnetic field, while many low-density striations seen outside of the filament extend along the magnetic field \\citep{palmeirim2013}. This geometry is sometimes believed to be an outcome of interstellar MHD (magnetohydrodynamic) turbulence \\citep{li2006}. That is, a turbulent sheet or filament is formed perpendicular to the global magnetic field.  There are a number of studies about the filamentary gas cloud based on the hydrostatic and magnetohydrostatic equilibria. When a filament is sufficiently long compared with its width, the cloud can be regarded as an infinitely long cylinder. \\revII{Under the assumptions of axisymmetry and no magnetic field,} density distribution of a cylindrical isothermal cloud with a central density $\\rho_c$ is expressed analytically \\citep{stodolkiewicz1963,ostriker1964} as \\be \\rho(r)=\\rho_c\\left(1+\\frac{r^2}{8H^2}\\right)^{-2} \\label{eqn:cyl-rho} \\ee where $H$ is a scale-height and is expressed using the isothermal sound speed $c_s$, the central density $\\rho_c$, and the gravitational constant $G$ as $H=c_s/(4\\pi G \\rho_c)^{1/2}$. This leads to a mass distribution, which is defined as a mass contained inside radius $r$ per unit length, such as \\begin{mathletters} \\begin{eqnarray} \\lambda(r) &=& \\int_0^r 2\\pi \\rho r dr\\\\ &=&\\frac{2c_s^2}{G}\\frac{{r^2}/{8H^2}}{1+{r^2}/{8H^2}}. \\label{eqn:cyl-lambda-dens} \\end{eqnarray} \\end{mathletters} \\hspace*{-3mm} The solution is truncated at a radius $R$ where the pressure $p=\\rho c_s^2$ is balanced with the external ambient pressure $p_{\\rm ext}$, that is, $\\rho \\ge p_{\\rm ext}/c_s^2$. Equation (\\ref{eqn:cyl-lambda-dens}) shows that the cylindrical filament has a maximum line-mass (mass per unit length) $\\lambda(R) \\le \\lambda(R=\\infty)=2c_s^2/G$, which corresponds to the mass of a filament in the vacuum $p_{\\rm ext}=0$. The character of the isothermal filament is controlled by a parameter as $\\lambda(R)/(2c_s^2/G)$ \\citep{nagasawa1987,inutsuka1992,fischera2012}. Herschel's observations of Aquila and Polaris clouds show that the Aquila main cloud has a large line-mass as $\\lambda \\gtrsim 5 \\times (2c_s^2/G)$ and is rich in protoclusters but that a portion with low line-mass as $\\lambda \\lesssim 2c_s^2/G$ is devoid of prestellar cores and protostars \\citep{andre2010} (To derive mass per unit length from observed column density, % the width of the filament is assumed constant in their paper as FWHM $\\sim 14000 {\\rm AU}$ in Aquila and $\\sim 9000{\\rm AU}$ in Polaris).  When the filaments have magnetic fields only parallel to their axes, $B_z$, the structure is also analytically given by equation (\\ref{eqn:cyl-rho}) but in this case $H=c_s(1+\\beta^{-1})^{1/2}/(4\\pi G \\rho_c)^{1/2}$, where the plasma beta is assumed to be constant as $\\beta=c_s^2\\rho/(B_z^2/8\\pi)={\\rm const}$ \\citep{stodolkiewicz1963}. The mass distribution $\\lambda(r)$ increases with the magnetic field strength as \\begin{equation} \\lambda(r) = \\frac{2c_s^2}{G}\\left(1+\\beta^{-1}\\right)\\frac{{r^2}/{8H^2}}{1+{r^2}/{8H^2}}. \\label{eqn:cyl-lambda-mag-dens} \\end{equation} Comparing equations (\\ref{eqn:cyl-lambda-dens}) and (\\ref{eqn:cyl-lambda-mag-dens}), the line-mass increases in proportion to $1+\\beta^{-1}$. The maximum line-mass supported against the self-gravity (for $r \\gg H$) increases also in proportion to $1+\\beta^{-1}$. Similar solutions are obtained numerically for the case whose mass-to-flux ratio is constant $\\Gamma_z\\equiv B_z/\\rho={\\rm const}$ \\citep{fiege2000a,fiege2000b}. In both cases, the poloidal magnetic field $B_z$ has the effect of increasing the mass of a static filament. \\citet{fiege2000a,fiege2000b} also consider the effect of toroidal magnetic field $B_\\phi$ assuming a type of flux conservation, $\\Gamma_\\phi\\equiv B_\\phi/(r\\rho)={\\rm const}$. In this case, the toroidal magnetic field $B_\\phi$ has the opposite effect of reducing the supported mass, because the $B_\\phi$ exerts the ``hoop stress'' to compress the filament in the radial direction. \\footnote{\\revII{The radial Lorentz force coming from $B_\\phi$ is proportional to the current in $z$-direction, $\\propto \\dif{rB_\\phi}{r}=\\dif{r^2\\rho\\Gamma}{r}$. Thus, the direction of the force is determined from the density distribution, or $M\\equiv \\dif{\\log\\rho}{\\log r}+2$. \\citet{fiege2000a} have shown that $M>0$ for their all isothermal and logatropic models, which means this force is working inwardly.}} However, the relation between the axis of the filament and the interstellar magnetic field is observed far from such simple configurations previously studied; that is, the actual filament is often perpendicular to the global magnetic field rather than the filament simply having poloidal and/or toroidal components. In the present paper, we revisit the magnetohydrostatic configuration of isothermal filaments, paying attention to the polarization observations, which indicate that the global magnetic field is often perpendicular to the filament.  The axisymmetric cloud threaded by the poloidal magnetic field has a similar maximum mass that depends on the magnetic flux. \\revII{From numerically obtained magnetohydrostatic configurations, it is shown that the maximum column density $\\sigma_{\\rm cr}$ depends on the magnetic flux density $B_0$ as $\\sigma_{\\rm cr}\\sim 0.17 B_0/G^{1/2}$(eq.(4.8) of \\citet{tomisaka1988b}).} This gives the maximum mass $M_{\\rm cr}$ is proportional to the magnetic flux $\\phi_0$, $M_{\\rm cr}\\sim 0.17 \\phi_0/G^{1/2}$. \\revII{This maximum column density is nearly equal to the maximum stable column density against the gravitational instability of magnetized plane-parallel sheet, $\\sigma_{\\rm cr} = B_0/(2\\pi G^{1/2})$ obtained by \\citet{nakano78}.}   The filamentary structure may be in dynamical contraction \\citep{inutsuka1992, kawachi1998} rather than in a hydrostatic state considered here. However the condition to begin the dynamical contraction is given from the hydrostatic maximum line-mass supported against the self-gravity.  The structure of this paper is as follows: in $\\S$2, the model and formulation for obtaining a magnetohydrostatic configuration are given. The method is a self-consistent field method similar to \\citet{mouschovias1976a,mouschovias1976b} and \\citet{tomisaka1988a,tomisaka1988b}, although these authors considered a disk-like cloud threaded perpendicularly by the magnetic field. Formulation for the filament with a lateral magnetic field is described in the following section. We show the numerical result in $\\S$ 3, in which structure of the filament is shown. Discussion on the effect of a magnetic field, such as how much mass is supported by the lateral magnetic field, is shown in $\\S$ 4. Section 5 is devoted to a summary. ",
        "conclusions": "\\subsection{Maximum Mass} \\label{sec:4.1} Figure \\ref{fig3} shows the masses of the equilibrium solutions. As is expected from the non-magnetized model, the line-mass is an increasing function of the central density or central-to-surface density ratio. The non-magnetized filament has a maximum line-mass of $\\lambda'_0=8\\pi$ (see eq.[\\ref{eqn:lambda-nondim}]), which corresponds to the dimensional value of $\\lambda_0=2c_s^2/G$ (see eq.[\\ref{eqn:lambda-dim}]), which is reached for $\\rho_c\\rightarrow \\infty$. Similarly to the non-magnetized model, the maximum line-mass for given parameters $R_0$ and $\\beta_0$ seems to be reached when $\\rho_c\\rightarrow \\infty$. The maximum line-mass should be estimated from the line-mass of the filament with finite central densities. We calculate $\\dif{\\log \\lambda_0}{\\log \\rho_c}$ for the state with the highest central density. When the slope $\\dif{\\log \\lambda_0}{\\log \\rho_c}$ is as small as $ < 0.1$, the maximum line-mass given from the largest line-mass obtained is a good estimation as the maximum line-mass that can be supported. The largest line-masses obtained for respective models are plotted in Figure \\ref{fig9} and shown in Table \\ref{table1}. The line-mass ($y$-axis) is plotted against the magnetic flux per unit length $\\Phi_{\\rm cl} \\equiv R_0B_0$ ($x$-axis). Asterisks represent the maximum line-mass for more reliable models, as   $\\dif{\\log \\lambda_0}{\\log \\rho_c} \\le 0.1$, while the cross represents the maximum line-mass for the models with $0.1 < \\dif{\\log \\lambda_0}{\\log \\rho_c} \\le 0.25$. From the asterisk points, we obtain an empirical formula for the maximum line-mass as \\be \\lambda'_{\\rm max} \\simeq 4.3 \\Phi_{\\rm cl}' + 20.8, \\label{eqn:-lambda'} \\ee using the least-squares method, where we add $'$ to emphasize that the quantities are normalized. Although the maximum line-mass of the non-magnetized filament given from the formula $\\lambda'_{\\rm max}(\\Phi_{\\rm cl}'=0)\\simeq 20.8$ is somewhat smaller than that obtained analytically $\\lambda'_{\\rm max}(\\Phi_{\\rm cl}'=0)=8\\pi\\simeq 25.1$, the fit is remarkable. % Since the line-mass and the magnetic line-flux are normalized by $c_s^2/4\\pi G$ and $c_s^2/(G/2)^{1/2}$ in this paper, the dimensional maximum line-mass is expressed with the dimensional flux as \\be \\lambda_{\\rm max} \\simeq 0.24 \\frac{\\Phi_{\\rm cl}}{G^{1/2}}+1.66 \\frac{c_s^2}{G}. \\label{eqn:mcr_nond} \\ee The critical mass-to-flux ratio has been obtained for the disk-like cloud as $(G^{1/2}M/\\Phi_B)_{\\rm crit}\\simeq 0.18 \\simeq 1/2\\pi$ (eq.(4.1) of \\citet{tomisaka1988b}; see also \\citet{nakano78}). It should be noticed that the factor for the filament, $0.24$, is similar to that of the disk, $0.18$\\footnote{ The mass-to-flux ratio ($M/\\phi$) and the line-mass to flux per unit length ($\\lambda/\\Phi_{\\rm cl}$) ratio have the identical dimension to the ratio between column density and magnetic flux density ($\\sigma/B$). Increasing the strength of magnetic field, a flat structure appears near the center of the cloud both in a disklike cloud \\citep{tomisaka1988b} and in a filamentary cloud. Thus, a flat structure threaded perpendicularly by the magnetic field appears commonly and such a disk seems to control the maximum mass and line-mass. If the condition of criticality is related to the structure of this flat part and thus related to the $\\sigma/B$ ratio of the structure, this explains the similarity of the factors appeared in the critical $M/\\phi$ and in the critical $\\lambda/\\Phi_{\\rm cl}$.}. Finally, the maximum line-mass (eq.[\\ref{eqn:mcr_nond}]) is finally evaluated as \\be \\lambda_{\\rm max} \\simeq 22.4  \\left(\\frac{R_0}{0.5\\,{\\rm pc}}\\right) \\left(\\frac{B_0}{10\\,{\\rm \\mu\\, G}}\\right) M_\\odot\\,{\\rm pc}^{-1} +13.9\\left(\\frac{c_s}{190\\, {\\rm m\\,s^{-1}}}\\right)^2 M_\\odot\\,{\\rm pc}^{-1}. \\label{eqn:mcr_dim} \\ee Thus, when the magnetic flux per unit length is larger than $\\Phi \\gtrsim 3\\, {\\rm pc \\,\\mu G}(c_s/190\\,{\\rm m\\,s^{-1}})^2$, it is concluded that the maximum line-mass of the filament is affected by the magnetic field (the first term is larger than the second term).  The factors in front of $\\Phi'_{\\rm cl}$ and $\\Phi_{\\rm cl}$ in equations (\\ref{eqn:-lambda'}) and (\\ref{eqn:mcr_nond}) must depend on the way of mass loading adopted here (eq.[\\ref{eqn:model_dm_dPhi}]). Analogy from the case of disk-like cloud (\\S IVb of \\citet{tomisaka1988b}), the factor may increase if we choose more uniform $d\\lambda/d\\Phi_{\\rm cl}$ rather than the centrally concentrated one assumed in equation (\\ref{eqn:model_dm_dPhi}).  \\subsection{Virial Analysis} In this section, we analyze the equilibrium state of the cloud using a virial analysis. The equation of motion for MHD is written \\be \\rho\\left[\\difd{\\bf u}{t} +\\left({\\bf u}\\cdot\\nabla\\right){\\bf u}\\right]=-\\nabla p -\\rho\\nabla\\psi +\\frac{1}{4\\pi}\\left(\\nabla\\times {\\bf B}\\right)\\times {\\bf B},\\label{eqn:eq-motion} \\ee where {\\bf u} represents the flow velocity. We consider an axisymmetric filament uniform in the axis direction. \\revII{Using the cylindrical coordinates, we multiply the position vector from the center, and integrate over the filament. The first term of the right-hand side gives} \\be \\int_0^R -\\difd{p}{r}r2\\pi r dr=-\\left[p 2 \\pi r^2\\right]_0^{R}+\\int_0^Rp4\\pi rdr =-2\\pi R^2p_s+2c_s^2\\lambda,\\label{eqn:virial1} \\ee where $R$, $p_s$, and $\\lambda$ represent the radius of the filament, the pressure at the surface of the filament, and the line-mass defined as $\\lambda=\\int_0^R \\rho 2\\pi r dr$. Treating the second term of the right-hand side of equation (\\ref{eqn:eq-motion}) in a similar way provides \\be \\int_0^R -\\rho \\difd{\\psi}{r}r 2 \\pi rdr =-\\int_0^R 2G\\lambda_r \\rho 2\\pi r dr =-2G\\int_0^R \\lambda_r d\\lambda_r =-G\\lambda^2, \\label{eqn:virial2} \\ee where we use the Poisson equation for the gravity as $ -\\dif{\\psi}{r}=-{2G\\lambda_r}/{r} $ and the definition of the line-mass contained inside the radius $r$ as $ \\lambda_r=\\int_0^r \\rho 2\\pi r dr. $ Thus, for a non-magnetic filament ${\\bf B}=0$, equation (\\ref{eqn:eq-motion}) gives the virial equation in equilibrium \\be 2\\pi R^2 p_s = G \\lambda\\left(\\frac{2c_s^2}{G}-\\lambda\\right) \\label{eqn:eq-of-B=0} \\ee Since the right-hand side of the equation must be positive, the line-mass $\\lambda$ has a maximum, $\\lambda_{\\rm max}=2c_s^2/G$. This critical line-mass is estimated as $\\lambda_{\\rm max}\\simeq 17\\, M_\\odot{\\rm pc^{-1}}(c_s/190\\,{\\rm m\\,s^{-1}})^2$ for the interstellar molecular gas with $T=10\\, {\\rm K}$.  When the filament is laterally threaded by the magnetic field, the last term of the right-hand side of equation (\\ref{eqn:eq-motion}) gives the magnetic force term and is estimated as \\be \\frac{B^2}{8\\pi}\\pi R^2 =\\frac{\\Phi_{\\rm cl}^2}{8}, \\label{eqn:magnetic-term} \\ee where $\\Phi_{\\rm cl}$ represents the magnetic flux per unit length as $\\Phi_{\\rm cl}=BR=B_0R_0$. Here we assumed that the flux $\\Phi_{\\rm cl}$ is \\revII{the same as in the parent state} in which the radius and the magnetic field strength are equal to $R_0$ and $B_0$, respectively. Equation (\\ref{eqn:eq-of-B=0}) becomes \\be 2\\pi R^2 p_s = 2c_s^2\\lambda-G\\lambda^2+\\frac{\\Phi_{\\rm cl}^2}{8}. \\label{eqn:eq-of-B<>0} \\ee Thus, the maximum line-mass increases and is equal to \\be \\lambda=\\frac{c_s^2+(c_s^4+G\\Phi_{\\rm cl}^2/8)^{1/2}}{G} \\sim \\frac{2c_s^2}{G}+\\frac{\\Phi_{\\rm cl}^2}{16c_s^2}, \\label{eqn:mass-formula-small_phi} \\ee where to derive an approximate formula \\revII{for} $\\lambda$ \\revII{in the weakly magnetized case} we assumed $2c_s^2/G\\gg \\Phi_{\\rm cl}^2/(16c_s^2)$ or $\\Phi_{\\rm cl} \\ll 3\\, {\\rm pc\\,\\mu G}(c_s/190\\, {\\rm m\\,s^{-1}})^{2}$. The extra line-mass allowed by the magnetic field is expected to be % \\be \\Delta \\lambda \\sim 2.5\\, M_\\odot\\,{\\rm pc}^{-1} (R_0/0.1\\,{\\rm pc})^2(B_0/10\\,{\\rm \\mu G})^2(c_s/190\\,{\\rm m\\,s^{-1}})^{-2}. \\ee Equation (\\ref{eqn:mass-formula-small_phi}) means that the maximum line-mass is expressed as $\\lambda'_{\\rm max}=a\\Phi_{\\rm cl}'^2 + b$ in the limit of small magnetic flux per unit length. Using our result for the five models with $0 \\le \\Phi_{\\rm cl}' =R'_0/\\beta_0^{1/2} < 5$ (Models A, B1, C3, C4 and the non-magnetized model), the maximum line-mass is well fitted by the following formula as \\be \\lambda'_{\\rm max}=0.82 \\Phi_{\\rm cl}'^2 +25, \\ee or in the dimensional form as \\be \\lambda_{\\rm max}=0.033 \\Phi_{\\rm cl}^2/c_s^2 + 2.0 c_s^2/G. \\label{eqn:empirical-lambda-weak-B} \\ee Although the numerical factor in front of $\\Phi_{\\rm cl}^2$ is twice as large as that expected from the virial analysis, our empirical fit [eq.(\\ref{eqn:empirical-lambda-weak-B})] seems to have a theoretical base.  In the case of $2c_s^2/G \\ll \\Phi_{\\rm cl}^2/(16c_s^2)$ or $\\Phi_{\\rm cl} \\gg 3\\, {\\rm pc\\,\\mu G}(c_s/190\\, {\\rm m\\,s^{-1}})^{2}$, equation (\\ref{eqn:eq-of-B<>0}) gives \\be \\lambda \\simeq \\frac{c_s^2}{G}+\\frac{\\Phi_{\\rm cl}}{2^{3/2}G^{1/2}}+\\frac{2^{1/2}c_s^4}{G^{3/2}\\Phi_{\\rm cl}}, \\ee and the mass increment due to the magnetic field (the second term) is expected to be proportional to the magnetic flux per unit length for a filament with large $\\Phi_{\\rm cl}$ as follows: \\be \\Delta \\lambda \\sim 30\\, M_\\odot\\,{\\rm pc}^{-1} (R_0/0.5\\, {\\rm pc})(B_0/10\\, {\\rm \\mu G}). \\ee This  reproduces the empirical mass formula obtained numerically (eq.~[\\ref{eqn:mcr_dim}]). When the filament is magnetically supported, the maximum line-mass is proportional to the magnetic flux per unit length $\\Phi_{\\rm cl}$. Therefore, the functional form of the empirical mass formula in $\\S$\\ref{sec:4.1} also has a theoretical meaning.   \\subsection{Column Density Distribution}  As for the modeling of the density profile of the filament, Plummer-like profiles are commonly adopted as \\citep{nutter2008,arzoumanian2011} \\be \\rho(r)=\\frac{\\rho_c}{[1+(r/R_f)^2]^{p/2}}, \\ee which has the analytic expression also for the column density as \\be \\sigma(r)=A\\frac{\\rho_cR_f}{[1+(r/R_f)^2]^{(p-1)/2}}, \\label{eqn:plummer-sigma} \\ee where $A$ is a numerical factor. The distribution is determined by the central density $\\rho_c$, core radius $R_f$, and the density slope parameter $p$ \\citep{nutter2008,arzoumanian2011}. However, there is a contradiction between theory and observation. That is, Herschel observations indicate $p \\sim 2$, although the isothermal hydrostatic filament has $p=4$ (eq.[\\ref{eqn:cyl-rho}]). This is sometimes considered as a consequence of dynamical contraction. Dynamical contraction of a filament whose pressure obeys the polytropic relation $p\\propto \\rho^\\gamma$ is expressed by a self-similar solution and has a slope such as $\\rho \\propto r^{-2/(2-\\gamma)}$ \\citep{kawachi1998}. And the infall speed expected from the solutions is reported consistent with observations \\citep{palmeirim2013}. On the other hand, from a standpoint that the filament is hydrostatic, some ideas to change the gas equation-of-state are proposed. Here we study the possibility that the column density distribution is changed by a magnetic effect. We study the column density distribution expected for the magnetized filament.  In Figure \\ref{fig10}, the column density integrated along the direction of the $y$-axis, $\\sigma(x)=\\int\\rho dy$,  and that from the $x$-axis,  $\\sigma(y)=\\int\\rho dx$, are plotted respectively in panels ({\\it a}) and ({\\it b}), in which the column density is calculated for the state with $\\rho_c=300$ of Model C3 (Fig.\\ref{fig5}{\\it c}). To mimic this kind of observation, an additional background column density is added to the magnetohydrostatic solution. We assume the additional column density of 0\\%, 1\\%, 3\\%, and 7\\% of the column density observed at the center, that is, $\\sigma(x=0)$ and $\\sigma(y=0)$. Four curves in this figure correspond to these different backgrounds. To estimate the parameters, we fit the distribution at three radii: at $r=0$, from $\\sigma(x=0)$ and $\\sigma(y=0)$ we calculate the column density at the center, $A\\rho_cR_f\\equiv \\sigma_0$, in equation (\\ref{eqn:plummer-sigma}); we calculate $r=r_{1/2}$ at which the equation gives a half of the central column density $\\sigma(r_{1/2})=\\sigma_0/2$; we also calculate $r=r_{1/10}$ at which the equation gives one-tenth  of the central column density, $\\sigma(r_{1/10})=\\sigma_0/10$. The values of $\\sigma_0$, $r_{1/2}$, and $r_{1/10}$ for this model are shown in Table \\ref{table2}. The parameter $r_{1/2}$ represents the size of the core and the parameter $r_{1/10}$ is related to the steepness of the density distribution in the envelope. Parameters to fit the numerical solution by the Plummer-like distribution are also shown in Table \\ref{table2}. Fitted Plummer-like distributions are also shown by dashed-dotted lines in Figure \\ref{fig10}. Although the $\\chi^2$-fitting might be more appropriate than this three-point fitting, Figure \\ref{fig10} shows that our fitting works well.  This figure shows that the parameter $p$ takes a value $p \\gtrsim 4$ when the additional background is low ($\\lesssim 3\\%$) or no background is added (0\\%). This means that the filaments have $p\\sim 4$ distribution, even if the magnetic field is taken into account. On the other hand, if we add relatively large background column density, $7\\%$ of $\\sigma_0$, a shallow slope in the envelope appears and the slope parameter is as small as $p\\lesssim 3$. This means that the slope parameter $p$ is highly dependent on the completeness of background subtraction from an observation. When the background subtraction is incomplete, a shallow slope in the envelope and a small $p$ parameter are expected, even if the observed power $p \\simeq 2$ may have other origin."
    },
    "1402/1402.3802_arXiv.txt": {
        "abstract": "{} {A hard X-ray shortage, implying the cooling of the corona, was observed during bursts of IGR J17473-272, 4U 1636-536, Aql X-1, and GS 1826-238. Apart from these four sources, we investigate here an atoll sample, in which the number of bursts for each source is larger than 5, to explore the possible additional hard X-ray shortage during {\\it Rossi X-ray timing explorer (RXTE)} era.} {According to the source catalog that shows type-I bursts, we analyzed all the available pointing observations of these sources carried out by the {\\it RXTE} proportional counter array (PCA). We grouped and combined the bursts according to their outburst states and searched for the possible hard X-ray shortage while bursting.} {We found that the island states of KS 1731-260 and 4U 1705-44 show a hard X-ray shortage at significant levels of 4.5 and 4.7 $\\sigma$ and a systematic time lag of $0.9 \\pm 2.1$ s and $2.5 \\pm 2.0$ s with respect to the soft X-rays, respectively. While in their banana branches and other sources, we did not find any consistent shortage. } {} ",
        "introduction": "Disc-accreting X-ray binaries (abbreviated as XRBs hereafter) are known to display distinct spectral states according to the variation of their accretion rate during the evolution of an outburst \\citep{Lewin1993,Remillard2006,Esin1997}. At the high state, the spectrum is usually dominated by thermal components, and at the low state, significant non-thermal hard X-rays can be observed, which is thought to originate from the corona or the jet. Although XRBs have been discovered decades ago, the corona formation process is still unclear. In theory, it can result from evaporation \\citep{Meyer1994,Esin1997,Liu2007,Frank2002} or magnetic re-connection \\citep{Zhang2000,Mayer2007}.   A hard X-ray shortage, implying the cooling of the corona, was observed during bursts of IGR J17473-272, 4U 1636-536, Aql X-1, and GS 1826-238 \\citep{Maccarone2003,Chen2012,Ji2013,Chen2013,Ji2014}. Apart from these four sources, we investigate an atoll sample here (Table 1, 2), in which the number of bursts for each source is larger than 5, to explore the possible additional hard X-ray shortage during the {\\it Rossi X-ray timing explorer (RXTE)}  era. In this sample, we pay extra attention to the sources that show complete color-color diagrams (CCDs; see Table 1). According to the CCDs, we can identify outburst states at a time around the occurrence of each burst and categorize the bursts into hard and soft states, respectively. Finally, we find two atoll sources  (4U 1705-44 and KS 1731-260) that showed corona cooling while bursting. In this paper, we focus on the results derived on these two sources.   The neutron star KS 1731-260 is a transient located near the Galactic center (l=1.07$^\\circ$, b=3.66$^\\circ$) with a mass less than 2.1 $M_\\odot$ and a radius R $\\leqslant$ 12.5 km \\citep{Ozel2012}. It was discovered in 1989 \\citep{Syunyaev1990} and remained active until transitioning to quiescence in early 2001. The distance was estimated to be $\\sim$7 kpc \\citep{Muno2000}. The neutron star 4U 1705-44 is a persistently bright XRB in the direction of the Galactic bulge. \\citet{Christian1997} derived a distance of 11 kpc from the peak flux of photospheric radius expansion (PRE) bursts. The kilohertz quasi-periodic oscillations were discovered using observations with {\\it RXTE} by \\citet{Ford1998}. The spectral analysis and fluorescent iron line have been carried out by \\citet{DiSalvo2005}.  Section 2 introduces the data and the observations we analyze. Section 3 describes the results of the corona cooling and their intrinsic timescale of reheating.The summary and discussion of these results are provided in Section 4.   \\begin{table*} \\tiny \\begin{center} \\caption{The atoll sources that show complete CCDs and have a burst number larger than 5. The columns denote the source name, number of bursts, averaged persistent flux at 40-50 keV and 3.6-18 keV, and peak flux of bursts at 3.6-18 keV for hard and soft states.} \\begin{tabular}{|c|c|c|c|c|c|c|c|c|} \\hline &\\multicolumn{4}{|c|}{hard state} & \\multicolumn{4}{|c|}{soft state} \\\\ \\hline source        & number    &   persistent flux at &  persistent flux at  &    peak flux at             &    number         &    persistent flux at      &  persistent flux at  &    peak flux at             \\\\ &           &    40-50 keV (cts/s) &  3.6-18 keV (cts/s)  &  3.6-18 keV (cts/s)         &                   &     40-50 keV (cts/s)      &  3.6-18 keV (cts/s)  &  3.6-18 keV (cts/s)         \\\\ \\hline 4U 1608-52    &    21     &   0.220 $\\pm$ 0.018  &   106.67  $\\pm$ 0.07   &  7065.73  $\\pm$  24.88    &     12            &     0.011 $\\pm$ 0.036      &  436.73  $\\pm$  0.13      &  8760.11  $\\pm$  40.66            \\\\ 4U 1702-429   &    6      &   0.155 $\\pm$ 0.011  &   90.47   $\\pm$ 0.05   &  5253.94  $\\pm$  33.79    &     3             &    -0.042 $\\pm$ 0.006      &  131.64  $\\pm$  0.03      &  3539.94  $\\pm$  11.05            \\\\ 4U 1728-34    &    39     &   0.245 $\\pm$ 0.007  &   218.72  $\\pm$ 0.05   &  6123.01  $\\pm$  18.98    &     17            &     0.033 $\\pm$ 0.011      &  342.89  $\\pm$  0.09      &  4204.80  $\\pm$  19.96            \\\\ EXO 0748-676  &    77     &  -0.003 $\\pm$ 0.005  &    21.09  $\\pm$ 0.01   &  856.74   $\\pm$  3.69     &     9             &    -0.058 $\\pm$ 0.015      &  21.16   $\\pm$  0.04      &  1164.84  $\\pm$  13.12            \\\\ KS 1731-260   &    16     &   0.337 $\\pm$ 0.008  &   126.21  $\\pm$ 0.05   &  2377.92  $\\pm$  12.48    &     1             &    -0.142 $\\pm$ 0.023      &  455.30  $\\pm$  0.22      &  3792.18  $\\pm$  68.01            \\\\ 4U 1705-44    &    30     &   0.247 $\\pm$ 0.008  &    93.71  $\\pm$ 0.04   &  1486.46  $\\pm$  7.71     &     14            &     0.003 $\\pm$ 0.012      &  124.57  $\\pm$  0.06      &  2099.39  $\\pm$  12.49        \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}    \\begin{table*}\\tiny \\begin{center} \\caption{The sources that have a burst number larger than 5 but show no complete CCDs. The columns denote the source name, number of bursts, averaged persistent flux at 40-50 keV and 3.6-18 keV, and peak flux of bursts at 3.6-18 keV.} \\scalebox{1.0}{ \\begin{tabular}{|c|c|c|c|c|} \\hline source        & number   &  persistent flux at  &    persistent flux at  &   peak flux at          \\\\ &          &   40-50 keV (cts/s)  &    3.6-18 keV (cts/s)  &  3.6-18 keV (cts/s)     \\\\ \\hline 4U 0513-401      &    11               &  -0.033 $\\pm$ 0.034      &     34.89       $\\pm$  0.06    &  1126.33  $\\pm$ 10.28  \\\\ x1735-444        &     8               &  -0.026 $\\pm$ 0.010      &     228.46      $\\pm$  0.09    &  2374.16  $\\pm$ 18.75  \\\\ 4U 1820-30       &    12               &   0.230 $\\pm$ 0.009      &     227.46      $\\pm$  0.06    &  3906.67  $\\pm$ 19.17 \\\\ HETEJ1900.1-2455 &    10               &   0.203 $\\pm$ 0.015      &      60.76      $\\pm$  0.05    &  5635.34  $\\pm$  27.54 \\\\ 1M 0836-425      &    15               &   0.273 $\\pm$ 0.011      &      79.82      $\\pm$  0.04    &  1030.44  $\\pm$  8.41  \\\\ EXO 1745-248     &    16               &   0.478 $\\pm$ 0.012      &      215.18     $\\pm$  0.08    &  1540.94  $\\pm$ 11.48   \\\\ 4U 1916-053      &    12               &   0.072 $\\pm$ 0.009      &      27.45      $\\pm$   0.02   &   2031.05  $\\pm$  13.24\\\\ IGR J17511-3057  &     9               &   0.234 $\\pm$ 0.008      &      45.88      $\\pm$   0.03   &   2945.85  $\\pm$  18.60 \\\\ SLX 1744-300     &     8               &   0.089 $\\pm$ 0.010      &      82.29      $\\pm$   0.04   &    779.84  $\\pm$  10.01 \\\\ SAX J1750.8-2900 &     6               &   0.109 $\\pm$ 0.019      &      126.83     $\\pm$   0.09   &    2965.07 $\\pm$  22.84 \\\\ XTE J1759-220    &     6               &   0.045 $\\pm$ 0.011      &      25.47      $\\pm$   0.03   &    695.87  $\\pm$  10.90  \\\\ \\hline \\end{tabular} } \\end{center} \\end{table*} ",
        "conclusions": "We report on the hard X-ray shortage during bursts in KS 1731-260 and 4U 1705-44, suggesting the corona can be cooled by the shower of seed photons from type-I bursts.  For all the six sources (the previous four and the two reported in this paper) so far showing corona cooling, the reheating timescale is $\\sim$ seconds. However, a similar phenomenon cannot be observed for other sources in Table 1 and 2.   \\subsection{Comparison between the sources showing the corona cooling} Here, we find two more atoll sources, which behaved as possible hard X-ray shortages. These, together with the previous four sources (Aql X-1, IGR J 17473-2721, 4U1636-536, and GS 1826-238) constitute six sources exhibiting cooling of the corona during type-I bursts. However, the evolution of the outbursts in these sources and the position in the CCDs when the bursts occurred are quite different. The bursts in IGR J17473-2721 are embedded in a long-lived preceding hard state of a single outburst \\citep{Chen2012}. The accretion rate in GS 1826-238 ranges from $5\\% {L}_{\\rm edd}$ to $9\\% {L}_{\\rm edd}$, which is remarkably stable. The luminosity of 4U 1636-536, KS 1731-260, and 4U 1705-44 is  $3\\%-16\\% {L}_{\\rm edd}$, $6\\%-38\\% {L}_{\\rm edd}$, and $1.1\\%-70\\% {L}_{\\rm edd}$, respectively. The bursts in Aql X-1 are mostly located in the failed outbursts (the outbursts without state transitions). This implies that the cooling of the corona prompted by the bursts may be independent of the detailed accretion evolution. In addition, it seems that the different duration of the bursts (for example, $\\sim$ 30s in 4U 1705-44 and $\\sim$ 100s in GS 1826-238), which suggest the different proportion of hydrogen in accretion materials, do not greatly affect the reheating process because of their similar time delays. This result seems to suggest that the physical process dominating the corona is similar in different sources.     \\subsection{Cooling and reheating} By assuming the equipartition between the thermal and magnetic energy ($ \\frac{3}{2}nkT \\sim  B^2/8\\pi$) or the balance condition ($c_{\\rm s}^2 \\rho \\sim B^2/8\\pi$) in typical ADAF models, the magnetic field in the corona is $\\sim 10^6$ G. Given that the relation of the cooling efficiency between Compton scattering and synchrotron radiation is $\\frac{P_{\\rm comp}}{P_{\\rm sy}} \\simeq \\frac{U_{\\rm ph}}{U_{\\rm B}}$, the corona cooling is dominated by the inverse Comptonization. Since an effective Compton cooling of a pure electron plasma takes $\\sim 10^{-6}$ s, the discovered time lag of about a few seconds should be intrinsic to the corona recovery process.  In theory, evaporation models give a timescale of days, and it is usually used to explain the behavior of the corona evolution in long-term outbursts. \\citet{Ji2014} proposed that the timescale for evaporation while bursting could drop significantly after considering the additional angular momentum transfer mechanism. The detailed reheating model, however, remains unknown.  Alternative to the evaporation model, the corona can also be formed through a magnetic re-connection process, as was firstly proposed according to the similarity between the XRB disks and the solar corona. In a magnetic re-connection model, similar to what happens in Sun, magnetic turbulence and buoyancy can trigger magnetic flares and cause an intense heating  \\citep{Zhang2000}. Following the model by \\citet{Liu2002}, the energy balance in the corona between heating by magnetic reconnection and cooling by Compton scattering, is as follows: \\begin{equation} \\label{ene_bal} \\frac{{{B^2}}} {{4\\pi }}{V_A} \\approx \\frac{{4kT}} {{{m_e}{c^2}}}n{\\sigma _T}c{U_{\\rm rad}l}, \\end{equation} where $V_A \\sim {B}/{\\sqrt{4 \\pi \\rho_0 }}$ is Alfv\\'{e}n speed. Here, $\\rho_0$ is the electron density, $\\sigma _T$ is the Thomson cross-section , $U_{\\rm rad}$ is the soft photon field to be Compton scattered,  $U_{\\rm rad} \\sim {\\sigma T_ {\\rm disc}^4}/{c} \\sim 2 \\times {10}^{13}$ erg {cm}$^{-3}$, $\\sigma$ is the Stefan-Boltzmann constant, and $l$ is the characteristic length of the magnetic loops. We estimate the timescale of the reheating process in a magnetic reconnection model as $l/c_{\\rm s}$, where $c_{\\rm s}$ is the sound speed and $c_{\\rm s} \\simeq 10 (T/10^4)^{1/2}$ km s$^{-1}$ \\citep{Frank2002}. By assuming the $l\\sim 10\\ r_g$ and a characteristic temperature between $10^7 K$ and $10^9 K$, the time delay ranges from 0.01 to 0.1s.   \\subsection{The atoll sources without hard X-ray shortage} Except for KS 1731-260 and 4U 1705-44, we find no significant hard X-ray shortage in other sources listed in both Table 1 and Table 2. During banana states for atoll sources enclosed in Table 1, there is no evidence of hard X-ray shortage probably because the persistent hard X-ray flux itself is too weak to be detectable. This could be the case as well for some sources listed in Table 2.   Apart from  KS 1731-260 and 4U 1705-44, there are seven other sources in Tables 1 and 2 that have the 40-50 keV flux larger than 0.2 cts/s. However, we see no obvious hard X-ray shortage while bursting from these seven sources. One possibility  that makes shortages invisible could be that, bursts with high temperature contribute hard X-rays that are not negligible while bursting to dilute the original shortage. For example, the mean temperature for 4U 1728-34 around the peak of the bursts is 2.8 keV, which leads to a contribution of 0.45 cts/s at 40-50 keV. As once discussed in Chen et al. (2012) in the case of the lacking shortage in the lagging low hard state of IGR J17473-2721, an additional possibility is that the corona may be located too far away from the neutron star to be effectively cooled off by the bursts. This occurs because in current observations the innermost radius of the accretion disk in the hard state can not be well constrained due to the relative low luminosity and severely hardened spectrum.  Another possibility could be that the hard X-rays are from the jet, which is usually hard to be cool off  under the soft X-ray shower of the bursts.    \\subsection*"
    },
    "1402/1402.4613_arXiv.txt": {
        "abstract": "Nonlocal massive gravity can provide an interesting explanation for the late-time cosmic acceleration, with a dark energy equation of state $w_{\\rm DE}$ smaller than $-1$ in the past. We derive the equations of linear cosmological perturbations to confront such models with the observations of large-scale structures. The effective gravitational coupling to nonrelativistic matter associated with galaxy clusterings is close to Newton's gravitational constant $G$ for a mass scale $m$ slightly smaller than today's Hubble parameter $H_0$. Taking into account the background expansion history as well as the evolution of matter perturbations $\\delta_m$, we test for these models with Type Ia Supernovae (SnIa) from  Union 2.1, the cosmic microwave background (CMB) measurements from Planck, a collection of baryon acoustic oscillations (BAO), and the growth rate data of $\\delta_m$. Using a higher value of $H_0$ derived from its direct measurement ($H_0 \\gtrsim 70$ km s$^{-1}$ Mpc$^{-1}$) the data strongly support the nonlocal massive gravity model ($-1.1 \\lesssim w_{\\rm DE} \\lesssim -1.04$ in the past) over the $\\Lambda$CDM model ($w_{\\rm DE}=-1$), whereas for a lower prior (67 km s$^{-1}$ Mpc$^{-1}$ $\\lesssim$ $H_0 \\lesssim 70$ km s$^{-1}$ Mpc$^{-1}$) the two models are statistically comparable. ",
        "introduction": "Modified gravitational theories have received much attention in connection to the dark energy problem \\cite{review}. In particular, the recent observational constraints derived from Planck and other data show that the dark energy equation of state $w_{\\rm DE}$ smaller than $-1$ is favored \\cite{Ade:2013zuv}. This may imply the infrared modification of gravity from general relativity (GR), because the models in the framework of GR --- such as quintessence \\cite{quin} and k-essence \\cite{kes} --- generally predict $w_{\\rm DE}$ larger than $-1$.  So far many dark energy models based on the large-distance modification of gravity have been proposed --- including the Dvali-Gabadadze-Porrati (DGP) model \\cite{DGP}, $f(R)$ gravity \\cite{fR}, Brans-Dicke theories \\cite{Brans}, and Galileons \\cite{Galileons}. In the DGP model the cosmic acceleration can be realized by the gravitational leakage to the fifth dimension, but it suffers from the incompatibility with observations \\cite{DGPcon} as well as the ghost problem \\cite{Luty}. In $f(R)$ gravity and Brans-Dicke theories it is possible to construct viable dark energy models at the expense of designing scalar potentials to be compatible with both cosmological and local gravity constraints \\cite{fR2}. In covariant Galileons there exists a tracker solution along which $w_{\\rm DE}$ evolves from $-2$ (matter era) to $-1$ (de Sitter era) \\cite{DT12}, but only the late-time tracking solutions are allowed from the joint data analysis of SnIa, CMB, and BAO \\cite{Nesseris}.  Besides the theories mentioned above, massive gravity has recently received significant attention due to the possibility of the late-time cosmic acceleration with a mass scale $m$ of the order of today's Hubble parameter $H_0$. In the original Fierz-Pauli theory \\cite{fierzpauli} there exists a so-called van Dam-Veltman-Zakharov (vDVZ) discontinuity \\cite{DVZ} with which the linearized GR cannot be recovered in the $m \\to 0$ limit. In the presence of nonlinear interactions the problem of the vDVZ discontinuity can be cured \\cite{Vainshtein}, but an instability mode called the Boulware-Deser (BD) ghost appears due to nonlinearities \\cite{BDghost}.  De Rham, Gabadadze and Tolley (dRGT) constructed a massive gravity theory \\cite{deRham:2010kj} in which the BD ghost is absent. In addition to acausality of the theory \\cite{aca} and the requirement of an external reference metric, there are some problems when the dRGT theory is applied to the cosmology. On the homogenous and isotropic cosmological background, it was shown that at least one ghost exists among five propagating degrees of freedom \\cite{DGM}. The self-accelerating solutions in the dRGT theory are also unstable against scalar and vector perturbations \\cite{ins}. The possible way out of these problems is to break the homogeneity or the isotropy of the Universe \\cite{homoge,ani} or to introduce other degrees of freedom \\cite{quasi1,quasi2,quasi3}.  An alternative approach to massive gravity was recently suggested by Jaccard {\\it et al.} \\cite{Jaccard}, who introduced nonlocal terms to obtain fully covariant equations of motion without referring to any external reference metric (see also Refs.~\\cite{DHK,Deser,Jhingan,Koivisto,Denonlocal,Zhang,Elizalde,Park,Deser13} for related works). This theory --- dubbed non-local massive gravity (NLMG)--- respects causality and reduces to a massless one without the vDVZ discontinuity in the $m \\to 0$ limit. The covariant equations of motion are given by \\be G_{\\mu\\nu}-m^2 (\\square_{\\rm ret}^{-1}G_{\\mu\\nu})^{\\rm T} =8\\pi G\\,T_{\\mu\\nu}\\,, \\label{nonlocal0} \\ee where $G_{\\mu \\nu}$ is the Einstein tensor, $T_{\\mu \\nu}$ is the energy-momentum tensor, $G$ is the gravitational constant, $\\square_{\\rm ret}^{-1}$ is the inverse of d'Alembertian operator computed with the retarded Green's function, and the superscript T represents the extraction of the transverse part.  The background cosmological dynamics based on Eq.~(\\ref{nonlocal0}) was studied in Ref.~\\cite{Modesto:2013jea}. There is a rapidly growing scalar mode responsible for the late-time cosmic acceleration, in which case the dark energy equation of state evolves from $w_{\\rm DE}=-1.725$ (matter era) to $w_{\\rm DE}=-1.506$ (accelerated era). Since the Planck data combined with the SnIa and WMAP polarization data placed the bound $w_{\\rm DE}=-1.13^{+0.13}_{-0.14}$ (95\\,\\%\\,C.L.) for constant $w_{\\rm DE}$ \\cite{Ade:2013zuv}, the NLMG model (\\ref{nonlocal0}) is in tension with the current observations of CMB and SnIa. In order to avoid the rapid growth of the scalar mode, we also require that the mass $m$ is much smaller than $H_0$.  Alternatively, Maggiore \\cite{Maggiore:2013mea} proposed a model given by the field equation \\be G_{\\mu\\nu}-\\frac13 m^2 (g_{\\mu \\nu}\\square_{\\rm ret}^{-1}R)^{\\rm T} =8\\pi G\\,T_{\\mu\\nu}\\,, \\label{nonlocalmo} \\ee where $g_{\\mu \\nu}$ is the metric tensor and $R$ is the Ricci scalar. In this case the strong instability of a scalar mode present in the theory (\\ref{nonlocal0}) is avoided, so that the dark energy equation of state does not significantly deviate from $-1$ ($w_{\\rm DE} \\approx -1.1$ in the deep matter era). The model has a predictive power due to the presence of a single parameter $m$ alone. For today's dark energy density parameter $\\Omega_{\\rm DE}^{(0)} \\simeq 0.68$, the mass $m$ is fixed to be $m \\simeq 0.67 H_0$ \\cite{Maggiore:2013mea,Foffacosmo}. It was also shown that the general relativistic behavior can be recovered inside the solar system \\cite{Keha}.  The retarded operator $\\square_{\\rm ret}^{-1}$ mentioned above is required for causality. However, the variation of some nonlocal action involving the inverse d'Alembertian operator $\\square^{-1}$ (such as $\\phi\\, \\square^{-1} \\phi$ with $\\phi$ being a scalar) usually symmetrizes the Green's function \\cite{Foffa:2013sma}. Hence the field equations of motion involving the retarded operator $\\square_{\\rm ret}^{-1}$ do not follow from a variational principle from some nonlocal action. Such a retarded nonlocal operator will only emerge when some classical or quantum averaging prescription is performed in a more fundamental local quantum field theory (QFT). In this sense, the field equations of motion (\\ref{nonlocal0}) and (\\ref{nonlocalmo}) should be considered only as effective classical equations that do not have a direct link to the action of a nonlocal QFT \\cite{Foffa:2013sma}. This is also the approach we follow in our paper.  Furthermore, if we use a quadratic action of gravitational waves associated with the perturbation equation of the theory (\\ref{nonlocalmo}) by simply replacing $\\square_{\\rm ret}^{-1}$ with $\\square^{-1}$, the resulting propagator apparently involves a ghostlike massive scalar \\cite{Maggiore:2013mea}. Foffa {\\it et al.} \\cite{Foffa:2013sma} showed that this apparent ghost is not a propagating degree of freedom and in the $m \\to 0$ limit it smoothly approaches a nonradiative degree of freedom of GR. This implies that we should regard Eq.~(\\ref{nonlocalmo}) as an effective classical equation of motion rather than promoting it directly to a full QFT (which typically involves some classical or quantum averaging). The issue of quantization---including ghosts---would be addressed in an underlying fundamental theory with a possible ultraviolet completion.  In this paper we study the cosmology (at the classical level) and observational constraints on the NLMG models. In Sec.~\\ref{backsec} the background equations of motion are derived for general models including (\\ref{nonlocal0}) and (\\ref{nonlocalmo}). We then discuss the evolution of $w_{\\rm DE}$ as well as the mass scale $m$ constrained from the background cosmology. In Sec.~\\ref{persec} we obtain the full equations of linear cosmological perturbations for the NLMG model (\\ref{nonlocalmo}). We also discuss the behavior of perturbations for the subhorizon modes relevant to large-scale structures. In Sec.~\\ref{resultscomp} we confront the NLMG model (\\ref{nonlocalmo}) with the latest observations of SnIa, CMB, BAO, and redshift-space distortions. Section ~\\ref{consec} is devoted to the conclusions. ",
        "conclusions": "\\label{consec}  In this paper we studied cosmological perturbations and observational constraints on the NLMG model. Our analysis of the background cosmology covers the two models given by Eqs.~(\\ref{nonlocal0}) and (\\ref{nonlocalmo}). We dealt with (\\ref{nonlocal1}) as an effective classical equation of motion for discussing the cosmology relevant to dark energy. The issues of ghosts and ultraviolet completion should be addressed in a more fundamental theory with a Lagrangian implementing quantum and classical averaging.  We derived the background equations of motion from (\\ref{nonlocal1}) on the flat FLRW background. For the models with $a_1 \\neq 0$ there is an instability for the field $V$ induced by the $-8V$ term on the lhs of Eq.~(\\ref{auto2}). In this case the mass $m$ is required to be much smaller than $H_0$ to avoid the early onset of the cosmic acceleration. Since the dark energy equation of state significantly deviates from $-1$, the models with $a_1 \\neq 0$ are strongly disfavored from the joint analysis of the SnIa, CMB, and BAO data.  For the models with $a_1=0$ the rhs of Eq.~(\\ref{auto2}) vanishes, so that the field $V$ does not grow for the appropriate initial conditions (\\ref{initial}). In order for the dark energy density $\\rho_{\\rm DE}$ to be positive, we found that the parameter $a_2$ has to satisfy the condition $a_2>0$. The dark energy equation of state $w_{\\rm DE}$ in the deep matter era can be estimated as Eq.~(\\ref{wdeana}), which shows good agreement with the numerically integrated solution ($-1.1<w_{\\rm DE}<-1.04$ for $-4<\\ln a<0$).  Expanding the field equations of motion and the metric to first order in perturbations about the flat FLRW background, we derived the full equations of cosmological perturbations for the NLMG model (\\ref{nonlocalmo}). The behavior of perturbations is also estimated for the modes relevant to galaxy clusterings. We found that the effective gravitational coupling $G_{\\rm eff}$ is very close to the gravitational constant $G$ for subhorizon perturbations characterized by the wave numbers (\\ref{modes}). Therefore, the evolution of $f\\sigma_8$ is similar to that in the $\\Lambda$CDM model (see Fig.~\\ref{fig3}). In this sense, the current growth-rate measurement alone is not able to distinguish between the NLMG and the $\\Lambda$CDM models.  We compared the NLMG model (\\ref{nonlocalmo}) against the latest cosmological observations, including SnIa, BAO, CMB, and redshift space distortions. Since the mass $m$ is not a free parameter, we developed an iterative algorithm to compute $m$ for each value of $\\Omega_{m}^{(0)}$ via Eq.~(\\ref{mH0}) and checked the consistency of Eqs.~(\\ref{auto1})---(\\ref{auto6}). The mass $m$ in the NLMG plays a similar role to the cosmological constant $\\Lambda$ in the $\\Lambda$CDM, such that the number of free parameters in the two models is the same. For the sake of comparison we considered the same parameters $(\\Omega_m^{(0)}, \\Omega_b^{(0)} h^2, \\sigma_{8,0})$ in both models.  The observational constraints on the NLMG model mainly come from the background expansion history rather than the growth history. The dark energy equation of state varies slowly from the deep matter era to today around $-1.1 \\lesssim w_{\\rm DE} \\lesssim -1.04$. Since $w_{\\rm DE}$ is approximately constant in the past, the situation is quite similar to the case of the constant $w_{\\rm DE}$ models studied in Ref.~\\cite{Shafer:2013pxa}. The likelihood results depend on the value of the $H_0$ prior due to the degeneracies of the CMB parameters.  We computed the chi squares in both the NLMG and the $\\Lambda$CDM models for several different values of $H_0$ ranging from the Planck best fit $h=0.673$ \\cite{Ade:2013zuv} to the Riess {\\it et al.} best fit $h=0.738$ \\cite{Riess:2011yx}. The results of our analysis are presented in Table \\ref{tab:results} and Figs.\\,\\ref{oms8plot}-\\ref{dchi2}. For $0.67 \\lesssim h \\lesssim 0.70$ the AIC shows that the NLMG and the $\\Lambda$CDM models are statistically comparable, but for $h \\gtrsim 0.70$ the NLMG model is strongly favored over the $\\Lambda$CDM model. We hope that future observations will pin down the values of $h$ to exquisite accuracy, clarifying whether the NLMG model is really preferred to the models with $w_{\\rm DE} \\ge -1$."
    },
    "1402/1402.3123_arXiv.txt": {
        "abstract": "{In view of interpreting the cosmological observations precisely, especially when they involve narrow light beams, it is crucial to understand how light propagates in our statistically homogeneous, clumpy, Universe. Among the various approaches to tackle this issue, Swiss-cheese models propose an inhomogeneous spacetime geometry which is an exact solution of Einstein's equation, while the Dyer-Roeder approximation deals with inhomogeneity in an effective way. In this article, we demonstrate that the distance-redshift relation of a certain class of Swiss-cheese models is the same as the one predicted by the Dyer-Roeder approach, at a well-controlled level of approximation. Both methods are therefore equivalent when applied to the interpretation of, e.g., supernova obervations. The proof relies on completely analytical arguments, and is illustrated by numerical results. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3409_arXiv.txt": {
        "abstract": "We present  a novel  method for detecting the relic neutrino background that takes advantage of  structured quantum degeneracy to amplify the drag force from neutrinos scattering off a detector.  Developing this idea, we present a characterization of the present day relic neutrino distribution in an arbitrary frame, including the influence of neutrino mass and neutrino reheating by $e^+e^-$ annihilation. We  present  explicitly the  neutrino velocity and de Broglie wavelength distributions for the case of an Earthbound observer.  Considering that relic neutrinos  could exhibit quantum liquid features at the present day temperature and density, we discuss the impact of neutrino fluid correlations on the possibility of resonant detection. ",
        "introduction": "Among the great science and technology challenges of this century is the development of the experimental capability to detect cosmic background neutrinos. To this end, we study the neutrino background as it is expected today, based on the freeze-out condition in the dense early Universe, combined with subsequent free-streaming and red-shifting to temperatures below the neutrino mass. In consideration of these results we argue that the present day massive relic neutrino background could form a quantum Fermi liquid state. We present arguments that the correlations inherent in a neutrino quantum Fermi liquid could be exploited in novel detection methods that utilize resonant amplification of relic neutrino drag forces.  This approach is in contrast to prior attempts to find a directly observable signature of the relic neutrinos, which have focused on the magnitude of the mechanical force due to scattering from a massless neutrino background~\\cite{Opher,Lewis,Opher2,Stodolsky:1975,Cabibbo:1982,Shvartsman,Langacker:1983,Smith,Ferreras:1995wf,Hagmann:1999kf,Duda:2001hd,Gelmini,Ringwald:2009,Liao:2012,Hedman}.  The consensus reached in the literature is that, for an unstructured (relativistic) neutrino gas, such effects are far from the abilities of current detector technology.  Aside from elastic scattering there are also inelastic processes -- we note   the development of the PTOLEMY experiment~\\cite{PTOLEMY} aiming  to observe relic electron neutrino capture by tritium, as originally proposed by Weinberg~\\cite{Weinberg}.  In this paper we will first characterize the free streaming distribution from the perspective of an observer in relative motion under the usual Boltzmann dilute gas assumption, utilizing the physically consistent equation of state from \\cite{Birrell:2013_2}.   We will then argue that high degree of degeneracy of the non-equilibrium relic neutrino distribution, together with their temperature $T_\\nu\\ll m_\\nu$, implies the inadequacy of the dilute gas assumption, resulting in a correlated background.  This leads us to explore the possibility of the detection of relic neutrinos by resonant amplification of the neutrino-detector interaction.  In our characterization of the cosmic neutrino background (CNB), we will focus on two important physical characteristics, namely the neutrino mass and effective number of  neutrinos $N_\\nu$. $N_\\nu\\ne 3$ signals either a yet undiscovered massless particle inventory in the Universe, or as we will follow in this work, a transfer of  $e^+e^-$ annihilation entropy into neutrinos resulting in  their momentum distribution being `hotter' than generally accepted~\\cite{Birrell:2013_2}. The magnitude of the  neutrino freeze-out temperature $T_k$ controls the amount of entropy that is transferred from $e^\\pm$ into neutrinos before freeze-out. % A numerical study based on the  Boltzmann equation  with two body standard model (SM) scattering~\\cite{Mangano:2005cc} gives $N_{\\nu}^{\\rm th}=3.046$.  $N_\\nu$ impacts Universe dynamics and recent  experimental  {\\small PLANCK CMB} results contain several fits~\\cite{Planck} which suggest that $N_\\nu\\simeq (3.30$--$3.62)\\pm0.25$. {\\small PLANCK CMB}  and lensing observations~\\cite{Battye:2013xqa} lead to  $N_{\\nu}=3.45\\pm0.23$.  In consideration of the complexity of low energy neutrino interactions with the dense $e^\\pm$ plasma near $T=O(1\\text{MeV})$, we treated the  freeze-out temperature $T_k$ as a free parameter  and obtained  the dependence of the free streaming neutrino distribution on $T_k$ in~\\cite{Birrell:2013_2}.  In particular, we find that the $T_k$ dependence can be expressed in terms of the measured  value  $N_\\nu$ and vice verse. This was done under the assumption of a strictly SM-particle inventory and allowed for a determination of the effect of neutrino mass $m_\\nu^i$, and  $N_\\nu $ on the equations of state determining the Universe expansion. A value of $N_{\\nu}\\simeq 3.5$  can be interpreted in terms of a delayed neutrino freeze-out during the $e^\\pm$ annihilation era. In the following we treat $N_\\nu$  as a variable model parameter within the general observed experimental range and use the above mentioned relations to characterize our results in terms of $N_\\nu$.  There are several available bounds on neutrino masses:\\begin{itemize} \\item[a)] Neutrino energy and pressure components are important before photon freeze-out and thus $m_\\nu$ impacts Universe dynamics. The analysis of cosmic microwave background (CMB) data alone leads to $\\sum_i m_\\nu^i<0.66$eV ($i=e,\\mu,\\tau$) and including Baryon Acoustic Oscillation gives $\\sum m_\\nu<0.23$eV~\\cite{Planck}.  {\\small PLANCK CMB} with lensing observations~\\cite{Battye:2013xqa} lead to  $\\sum m_{\\nu}=0.32\\pm0.081$ eV. \\item[b)] Upper bounds have been placed on the electron neutrino mass in direct laboratory measurements  $m_{\\bar\\nu_e}<2.05$eV~\\cite{PDG,Aseev}.\\end{itemize} In the subsequent analysis we will focus on the neutrino mass range $0.05$eV to $2$eV in order to show that direct measurement sensitivity allows for the exploration of a wide mass range.    In section \\ref{sec:dist_background} we introduce the free-streaming massive neutrino distribution.  We present our characterization of the neutrino distribution, including the dependence on $N_\\nu$ and $m_\\nu$, in section \\ref{sec:dist}. In section \\ref{sec:q_fluid} we show that at present day temperatures and densities, the CNB will have properties of a quantum Fermi fluid, including  density correlations. With this, we discuss the possibility of detection via resonance with the neutrino dynamical correlation frequency, which we estimate in the Earth frame to be within the range 110-160 MHz. ",
        "conclusions": "} In this work we have characterized the relic cosmic  neutrino spectra in terms of their velocity, energy, and de Broglie wavelength distributions in a frame of reference moving relative to the neutrino background, with all examples focused on an Earth bound observer. We have shown explicitly the mass, $m_\\nu$, dependence and the dependence on neutrino reheating expressed by $N_\\nu$, choosing a range within the experimental constraints.  The dependence on  $N_\\nu$ and $m_\\nu$ shown as a ratio on  linear scale in figure \\ref{fig:deBroglie_300} (bottom frame) is sufficiently strong to suggest  that if and when relic neutrino detection becomes possible both  $N_\\nu$ and $m_\\nu$ would be directly measurable. The effect of $N_\\nu$ as presented here is to increase neutrino flux~\\cite{Birrell:2013_2}, see \\req{nnu}. However, to this end one must gain precise control over the enhancement of neutrino galactic relic density due to  gravitational effects~\\cite{Ringwald:2004np} as well as the annual modulation~\\cite{Safdi}.  As we argued in section \\ref{sec:q_fluid}, the  relic  background of massive neutrinos will not behave as a purely free streaming gas, but have properties of a quantum liquid.  This inevitably leads to  impact correlations which will cause the power spectrum due to neutrinos  scattering off a detector to deviate from white noise, as we discuss in section \\ref{sec:q_fluid_detect}. The spectrum should exhibit a peak near a frequency  arising from the relative velocity of the terrestrial observer  with respect to CNB. Based on a simple model we find an expected range of  111 -- 140MHz for  the central value, with the frequency increasing based on gravitational accretion enhanced neutrino density. The frequency is subject to modulation by the temporal variation of the observer velocity vector with respect to CNB.   Such correlated structure in the noise spectrum opens the possibility for resonance based CNB detection methods, amplifying the otherwise negligible effect of neutrino drag."
    },
    "1402/1402.1968.txt": {
        "abstract": "Be stars exhibit variability for a great number of observables. Putting the pieces together of the disk dynamics is not an easy task and requires arduous modeling before achieving a good fit of the observational data. In order to guide the modeling process and make it more efficient, it is very instructive to investigate reference dynamical cases. This paper focuses on continuum polarimetric quantities and is the second of a series that aims to demonstrate the capacity of deriving the dynamical history and fundamental parameters of a classical Be star through the follow-up of various observables. After a detailed study of the different opacities at play in the formation of polarized spectra, we investigate predictions of polarimetric observables in the continuum for different dynamical scenarios. Our models are based on a coupling of a hydrodynamic viscous decretion simulations in a disk and a 3-D non-LTE radiative transfer code. Introducing the polarization color diagram (PCD), we show that certain combinations of polarimetric observables exhibit features that are characteristic of a mass loss history. This diagram also enables to estimate fundamental parameters such as the inclination angle, the disk density scale and the $\\alpha$ viscous diffusion parameter. We present the PCD as a powerful diagnosis tool to track the dynamical phases of a Be star such as disk building-up, dissipation, periodic and episodic outbursts. Finally we confront our models with observations of 4 Be stars that exhibited long-term polarimetric activity.  %Which band, niveau de polarization attendu, quels diagrammes sont caracteristiques de quels scenarios, alpha ??  %   comparaisons aux vrais donnees : ca se rapporte a des combinaisons de cas classiques ? ",
        "introduction": "Be stars are non-supergiant, early-type stars with a circumstellar (CS) disk that is created from matter ejected from the star. Recent observational facts brought by spectro-interferometry and spectro-astrometry \\citep[e.g.][]{2012A&A...538A.110M,2012MNRAS.423L..11W} supports the fact that so far all studied Be star disks rotate in a Keplerian fashion. This characteristic, together with other observational signatures of the disk outlined in \\cite{2011IAUS..272..325C}, are properties that only the viscous decretion disk  (VDD) model can reproduce. This model, first suggested by \\cite{1991MNRAS.250..432L} and further developed by \\citet{1997LNP...497..239B}, \\citet{1999A&A...348..512P}, \\citet{2001PASJ...53..119O}, \\citet{2005ASPC..337...75B} and \\citet{2008MNRAS.386.1922J}, among others, uses the angular momentum transport by turbulent viscosity to lift material into higher orbits, thereby causing the disk to grow in size. This model has already been successfully applied to systems showing stable continuum emission: e.g. $\\zeta$ Tauri \\citep{car09}, $\\chi$\\,Oph \\citep{2008ApJ...689..461T} and $\\beta$\\,CMi \\citep{2012MNRAS.423L..11W} and systems exhibiting a more variable photometric activity \\citep[28\\,CMa,][]{2012ApJ...744L..15C}.  In a recent review paper, \\cite{rev2013} discuss in detail the observational and theoretical evidences in support of the VDD scenario for Be stars.  Polarization is a powerful tool to study the geometry of the disk (opening angle, flaring) of Be stars without angularly resolving it. Polarized flux originates from electron scattering off the disk, and is affected by both pre- and post-scattering absorption by \\ion{H}{1} atoms \\citep{wood96b,2013ApJ...765...17H}. Since \\ion{H}{1} opacity depends on the physical state of the gas, studying polarimetric observables at different wavelengths allows one to probe different regions of the disk. In the literature, the polarimetric technique has an established history of providing a unique diagnosis in identifying and studying the detailed CS environments of Be stars \\citep[][]{1997ApJ...477..926W,2007ApJ...671L..49C,car09,2010ApJ...709.1306W,2011ApJ...728L..40D}.  Haubois et al. (2012), hereafter Paper I, studied the temporal variability of Be disks, based on {\\sc singlebe} VDD hydrodynamics simulations  \\citep{2007ASPC..361..230O}. {\\sc singlebe} solves the 1-D surface density evolution equation for a viscous isothermal Keplerian decretion disk. The effects of variable mass injection rates on the disk structure, and their corresponding effect on the photometry, were studied at different wavelengths and compared to observations. More specifically, we first studied the different timescales that characterize the evolution of the disk surface density and how this surface density responds to changes in the mass injection rates. These surface density profiles were used as inputs to the three-dimensional non-LTE Monte Carlo radiative transfer code {\\sc hdust}  \\citep{car06,2008ApJ...684.1374C} that allowed the calculation of photometric observables at various wavelengths. The characteristic shapes of these lightcurves agree qualitatively well with observations, which provide strong circumstantial evidence that viscosity is indeed the mechanism that redistributes matter along the CS disk. The first successful confrontation of theoretical VDD lightcurves with observations was done by \\cite{2012ApJ...744L..15C} for the Be star 28\\,CMa. To summarize, Paper I  provides the reader with a description of the photometric variability from a Be star in the framework of the VDD model. With the present paper, we aim at  exploring the variability of the  \\emph{continuum} polarimetric features in the same manner.  %We then computed lightcurves for Be stars subject to different mass loss rate histories and studied how the lightcurves depended on the other system parameters.  In \\S\u00f7\\ref{origin}, we describe the polarigenic mechanisms operating in Be disks. Then we present the dynamical models we investigated and their signatures on common continuum polarimetric observables. We also show the diagnosis potential of a series of diagrams that we named polarization color diagrams (PCDs) that represent a powerful tool to follow the mass injection history in Be stars (\\S\u00f7\\ref{prediction}). Finally, a discussion and a comparison to observed data of these synthetic observables are presented in \\S\u00f7\\ref{comp} before concluding.  %mass injection = (mass ejected from the central star and injected into the disk) ",
        "conclusions": "In this paper, we first review the mechanisms at the origin of the continuum polarization levels observed for Be stars under the assumption of a steady-state VDD model. For different disk densities and spectral types, we detail how the electron scattering, bound-free and free-free opacities shape the polarized spectrum. Analyzing further these opacities, we describe how they vary radially to establish the regions of the disk that are electron-scattering dominated and the regions that are bound-free dominated. For low densities (typically $1 \\times 10^{-12} \\rm g\\,cm^{-3}$), all models are electron scattering dominated and the polarization spectrum is roughly flat regardless the spectral type. For larger densities (typically greater than $1 \\times 10^{-11} \\rm g\\,cm^{-3}$), the bound-free opacity becomes the dominant opacity in the disk, at least in its innermost part. Based on this opacity description, we can define two polarimetric features that depend on different disk properties: $P_V$, the polarization level in the visible (or any other wavelength), is a measure of the scattering mass of the disk; and $P_{\\rm BD}$ (the polarization change across the Balmer jump), a measure of the \\emph{ color} of the polarization spectrum, which depends on the relative importance of the bound-free opacity to the total opacity. With a more realistic modeling based on a coupling of hydrodynamics and radiative transfer simulations, we show how these polarimetric features evolve with time for different disk mass injection scenarios. Different mass injection histories (constant, periodic or episodic) result in specific behaviors of the polarimetric observables. This led us to introduce an extension of the concept of BJV diagrams \\citep{2011ApJ...728L..40D} that we named PCD (Polarization Color Diagram) as an analogy to color magnitude diagrams. The PCD plots $P_{\\rm BD}$ (or other measure of the polarization color) versus the polarization level. PCDs constitute a powerful diagnosis tool to derive different physical parameters of the disk such as the inclination angle, the viscous coefficient $\\alpha$, the disk base density and the spectral type of the central star on top of the mass injection history. Typically, the polarimetric features of a star evolving from a diskless phase (normal B star) to a Be phase (with a disk present) and back to a B phase will appear as a loop in the PCD, the upper part of the loop being associated with the disk construction and the lower part with the dissipation. In some particular cases of very high densities and high inclination angles, morphologic variations of the PCDs with the apparition of a secondary loop were found. Albeit the useful diagnostic potential, the morphologies seen in the PCD bear some degeneracy with respect to some parameters. For this reason, we stress that the more observables and the better time coverage, the easier to infer a dynamical scenario and physical parameters of the system. %For instance, similar PCD loops can be obtained for a system with a rather low inclination angle ($\\sim 30\\degr$) and an early spectral type (B2) and for a system with a higher inclination  ($\\sim 70\\degr$) and a later spectral type (B5). The theoretical predictions were confronted with observed PCDs of four stars. While $\\pi$~Aqr exhibits quite a clear building-up/dissipation pattern in its PCD,  60 Cyg presents a more irregular case that is likely to be the result of mass injection rate variations. We then analyse the PCD of  $\\delta$ Sco and speculate that it probably represents the polarimetric counterpart of an outburst seen with photometry. Finally, we point out that the huge variations and the counter clockwise structure of $\\psi$ Per's PCD can be explained with our predictions for a B5 spectral type star, with a high inclination angle and a high density. However, a detailed modeling remain to be carried out for each specific star  to achieve an accurate reproduction of the observed PCDs.           %Viscous diffusion is improving, a non isothermal will be soon released. %Link toward the website : fix this color scale pbm wit IDL. %Further scenarios to investigate : more complex mass loss decretion, blob  puffing. %Opening on interferometry and spectroscopy."
    },
    "1402/1402.4703_arXiv.txt": {
        "abstract": "We present numerical simulations of active region scale flux emergence covering a time span of up to 6 days. Flux emergence is driven by a bottom boundary condition that advects a semi-torus of magnetic field with $1.7 \\times 10^{22}$~Mx flux into the computational domain. The simulations show that, even in the absense of twist, the magnetic flux is able the rise through the upper $15.5$ Mm of the convection zone and emerge into the photosphere to form spots. We find that spot formation is sensitive to the persistence of upflows at the bottom boundary footpoints, i.e. a continuing upflow would prevent spot formation. In addition, the presence of a torus-aligned flow (such flow into the retrograde direction is expected from angular momentum conservation during the rise of flux ropes through the convection zone) leads to a significant asymmetry between the pair of spots, with the spot corresponding to the leading spot on the Sun being more axisymmetric and coherent, but also forming with a delay relative to the following spot. The spot formation phase transitions directly into a decay phase. Subsurface flows fragment the magnetic field and lead to intrusions of almost field free plasma underneath the photosphere. When such intrusions reach photospheric layers, the spot fragments. The time scale for spot decay is comparable to the longest convective time scales present in the simulation domain. We find that the dispersal of flux from a simulated spot in the first two days of the decay phase is consistent with self-similar decay by turbulent diffusion. ",
        "introduction": "Active Regions (ARs) are prominent centerpieces of solar magnetic activity. In themselves, ARs exhibit interesting manifestations of the complex interplay between magnetohydrodynamics  (MHD) and radiative transfer. As demonstrated by numerical simulations in recent years, this interplay is responsible for a wide range of phenomenological features associated with ARs, including sunspot umbrae~\\citep{Schuessler:UmbralConvection,Rempel:SunspotStructure,Rempel:SunspotSubsurfaceStructure} and penumbrae~\\citep{Heinemann:PenumbraFineStructure,Rempel:SunspotStructure,Rempel:PenumbralStructure,Rempel:PenumbralFineStructure,Kitiashvili:SeaSerpentPenumbra,Rempel:LivingReviews,Rempel:Robustness}, light bridges~\\citep{Cheung:ARFormation}, pores~\\citep{Cameron:Pores,Kitiashvili:Spontaneous,SteinNordlund:ARFormation}, plages~\\citet{Voegler:MURaM} as well as anomalous features in solar granulation associated with emerging flux~\\citep{Cheung:FluxEmergenceInGranularConvection,Cheung:SolarSurfaceEmergingFluxRegions,MartinezSykora:TwistedFluxEmergence,TortosaAndreu:FluxEmergence,Fang:DynamicCoupling}.  Even before 3D numerical simulations addressing AR formation were feasible, it was generally accepted that ARs form when somewhat coherent, buoyant bundles of these toroidal fields rise toward the surface~\\citep[see reviews by][and references therein]{Moreno-Insertis:Review,Fisher:SolarDynamoAndEmergingFlux,Fan:LivingReview,Fan:LivingReview2009,Archontis:ReviewArticle}. The physical mechanisms that influence AR tilts and asymmetries have been addressed in idealized models that use the thin flux tube or analestic approximations. The assumptions employed in these models are not appropriate for the near-surface layers of the convection zone ($z \\lesssim -20$ Mm) and 3D, fully-compressible MHD simulations are required to investigate how ARs form at whitthe surface.  Radiative MHD simulations of AR formation following flux emergence by~\\citet{Cheung:ARFormation} and~\\citet{SteinNordlund:ARFormation} have partially bridged this gap in our models of the life cycle of ARs but outstanding questions remain. In both models, the magnetic field is given a chosen strength and orientation at the bottom boundary and the ensuing evolution is dependent on the choice. In the model of~\\citet{SteinNordlund:ARFormation}, horizontal magnetic field of constant strength and orientation is fed into the domain through the bottom boundary in upflow regions. In the case of~\\citet{Cheung:ARFormation}, the magnetic configuration was chosen to be a semi-torus to mimic the rise of a somewhat coherent flux loop from layers below the computational domain. In both cases,  the origin of the magnetic fields that are fed into the computational domains is not addressed, although the existence of buoyant magnetic flux tubes coherent over length scales exceeding tens of megameters is found to be self-consistently generated in recent global dynamo simulations~\\citep{Nelson:DynamoLoops}.  In the present work, we extend the approach taken in~\\citet{Cheung:ARFormation} to examine how buoyant flux tubes would emerge and evolve should they be present in the convection zone. In this article, we present numerical experiments based on radiative MHD that address certain aspects of the formation and decay of ARs. In one of the numerical experiments, we follow the full life cycle of the modeled AR (i.e. from pre-emergence to decay) over the course of $6$ days. We complement this simulation with control experiments that examine the role of the bottom boundary condition we use to initiate flux emergence. The combined results allow us to examine the robustness of the AR formation process and to examine the influences that lead to AR fragmentation. In particular we aim to address the following key questions: 1. How does magnetic field decouple from the enclosed mass to allow for flux emergence and spot formation in a highly stratified medium? 2. What is the role of subsurface flows in AR evolution? 3. What are the photospheric signatures of field aligned flows that are expected to occur as consequence of angular momentum conservation? 4. Which processes govern the decay of flux concentrations after flux emergence?   The remainder of the article is structured as follows. Section~\\ref{sect:setup} presents the setup of the numerical simulations. Section~\\ref{sect:results} presents results and lessons learned from the numerical simulations. Finally, section~\\ref{sect:discussion} discusses the implications of our findings in the context of emerging flux regions, AR formation and the solar dynamo. ",
        "conclusions": "\\label{sect:discussion} We presented a series of flux emergence simulations similar to \\citet{Cheung:ARFormation}. The main differences are a larger domain with about 3 times more overall flux and an about $10$ larger overall density contrast. In addition we considered here also a setup with a field aligned flow directed in the negative $x$-direction. Our simulations lead to the formation of a pair of spots with about $10^{22}$~Mx flux, which is about $50-60\\%$ of the flux we moved into the domain across the bottom boundary. Interestingly this ratio is not very different from \\citet{Cheung:ARFormation} despite the fact that we a) considered an untwisted flux tube, b) had $10$ times more density stratification in our domain, and c) started with an initially 2 times weaker magnetic field. Overall this points to some robustness of the processes we find during flux emergence and sunspot formation.  The addition of a field aligned flow, which is expected to arise as a consequence of angular momentum conservation, leads to two distinct photospheric signatures: 1. We find a significant asymmetry between the two spots. The spot that would correspond on the Sun to the leading spot is more coherent and axisymmetric. 2. We find systematic differences in the timing of the formation of both spots. The spot corresponding to the leading spot on the sun forms overall faster after an initial delay. While both spots reach about half of their peak flux at the same time, the formation of the left spot (corresponding to the trailing spot) starts earlier and continues past the formation of the right spot. This appears to be opposite to the observed behavior of the majority of ARs, although \\citet{McIntosh:BirthAndEvolutionOfSunspots} reports on 2 cases (out of 15) in which the trailing spot formation preceded the leading spot formation by $7$ and $26$ hours (we find about $10$ hours difference). In a recent observational study of the birth of two ARs with SDO/HMI data, the following spot seems to form earlier in one case (AR 11105) but later in the other~\\citep[AR 11211, see Figure in][]{Centeno:NakedAR}. While our simulation may not represent the \"typical\" case, it does not strictly contradict observations in this aspect either. Comprehensive statistical studies surveying spot formation times and other observables in leading and following polarities will be crucial to settle the question of what is a typical AR and what are the expected deviations from the typical AR. We note that our setup was chosen to highlight the influence of a field aligned flow in separation, rising flux tube simulations typically lead to additional asymmetries between both polarities that could potentially alter that outcome. We also note that the horizontal extent of the simulation domain in conjunction with periodic boundaries prevents the spread of magnetic field in the photosphere and potentially influences the timing of the left spot formation.  In our setup the total mass enclosed by the semi-torus we emerge across the bottom boundary corresponds to about $50\\%$ of the total mass in the domain. Consequently continuing flux emergence and spot formation is only possible if mass and magnetic field decouple on average, allowing for the mass to leave the domain through the open bottom boundary while more than $50\\%$ of the emerging flux reaches the top boundary of the domain (the total density contrast in the mean stratification is more than $10^6$). The decoupling of mass and magnetic field happens mostly through two processes. In the early stages of flux emergence persistent inflows at the bottom boundary condition lead to the formation of a subsurface field configuration with strong alignment between field and flow. This is mostly the case for the right spot with the stronger inflow at the footpoint. Due to flux freezing the overturning mass flux initially holds down most of the magnetic field and prevents emergence of flux into the photosphere (less than $0.1\\%$~of the mass flux present at the bottom boundary reaches the photosphere due to strong stratification). Since most flows are field aligned, the strong horizontal divergence of the velocity field does not weaken the field and leads to the built up of a several kG strong horizontal magnetic layer in the middle of the domain. At later stages the horizontal field rises toward the surface and emerges through the photosphere in form of granulation scale $\\Omega$-loops that allow for an effective decoupling of mass and magnetic field. Most of the magnetic flux emerges through the photosphere in form of mixed polarity field, spots form after a separation of polarities.  For both spots the field re-amplification in the photosphere is related to strong downflows, although the origin of these downflows differs. The left spot forms above the downflow region at the periphery of the flux emergence region. This downflow region is a consequence of mass conservation, i.e. the mass which is pushed into the domain through our emergence boundary condition has to leave the domain as we keep the average gas pressure at the bottom boundary unchanged. Since this downflow region shows an orientation in the y-direction, the initial shape of the left spot is also more sheet-like than axisymmetric. In the case of the right spot the longer lasting upflow prevents the drainage of low entropy material continuously forming in the photosphere. As a consequence we find and enhancement of the average upflow/downflow temperature contrast in about $8-10$~Mm depth beneath the photosphere. This effect is strongest in the center, where horizontally diverging flows are weakest. After the inflow decays away (as a consequence of switching the bottom boundary condition back to an open boundary), the resulting top heavy stratification leads to strong downflows, while the surrounding system of strong horizontal field is moving upward. Flux emerges in the photosphere in form of $\\Omega$-loops while the positive polarity field is amplified in downflows preferentially located above the foot-point of the right spot. The almost circular shape of the right spot is a reflection of the circular shape of preceding subsurface flow field with strong horizontal divergence.  In terms of the total unsigned flux the simulated active region transitions directly from formation to decay phase. For the right spot we find a short $\\sim 30$~hours stationary phase in terms of signed flux between formation and decay. This behavior is qualitatively similar to observational reports studying the formation and decay of ARs~\\citep[e.g.][]{McIntosh:BirthAndEvolutionOfSunspots,Verma:DecayingSunspot}, but we find substantial differences in the associated time scales. The decay is mostly driven by the largest scale convective motions present in our domain, i.e. the convection pattern near the bottom boundary. This sets both the decay time scale as well as surface magnetic field pattern during the decay phase. The time scale of decay is of the order of a few days, which is also consistent with the single spot simulations of \\citet{Rempel:SunspotSubsurfaceStructure}, but significantly shorter than the typical life time of sunspots. This reinforces their conclusion that realistic spot life times require simulations in deeper domains, likely around $30-50$~Mm depth. The photospheric magnetic field pattern during the decay phase approximately reflects the convective cell structure near the bottom boundary. In our setup it happened that new upflow cells formed at the bottom boundary beneath both spots leading to a ring-like appearance during the decay phase. This is a detail of the simulations that depends strongly on the location of the bottom boundary and would likely not happen in a substantially deeper domain.  In the initial stages of spot decay, the spreading of vertical flux over the photospheric surface was shown to be reasonably well-described in terms of a turbulent diffusion process. A self-similar solution to the diffusion equation for a Gaussian profile~\\citep[][]{Meyer:GrowthAndDecayOfSunspots,Mosher:PhD}with constant turbulent diffusivity was found to approximate the lateral migration of flux surfaces over the course of two days. During this time, the associated flux loss rates (from both the numerical simulation and analytical model) are in accordance with measured flux loss rates from observed sunspots as reported by~\\citet{Kubo:MMFs}. Thereafter, fragmentation of the spot led to increased rates of magnetic flux loss.  During formation and decay phase flux transport in the photosphere cannot be easily represented by simple expressions. We find that in general contributions from mean and fluctuating parts partially offset each other, similarly the contributions from radial and vertical expressions tend to have opposing signs. In particular the expression $\\overline{-v_rB_z}$ alone is insufficient to account for the flux transport rate and the complementary component $\\overline{v_zB_r}$ must also be taken into account.  The patterns of spot formation and decay we found in this investigation are a consequence of the flux emergence boundary condition we implemented. This raises the question which aspects can be considered robust and which aspects are very specific to this setup. Nevertheless, a comparison to \\citet{Cheung:ARFormation} already shows that neither twist, overall flux, field strength nor density contrast (domain depth) have a dramatic influence on the overall evolution. We have conducted additional control experiments using the same domain size and overall flux. We did not find a significant difference between simulations starting with $10$ and $20$~kG initial field strength. We also explored different bottom boundary conditions after the initial flux emergence. There is no significant difference between keeping the magnetic field fixed or allowing the foot-points to erode. The one aspect which is crucial is a decay of the initial upflow. If the upflow is continued at the bottom boundary no spot formation is observed. Since flux emergence is (observationally) a event with limited duration, using a boundary condition that switches after about $20$ hours of emergence to an open boundary as we did seems very reasonable. Identifying the underlying physical process is beyond the scope of our numerical setup, but will be addressed in the future as larger and deeper domains become available.  Recently \\citet{SteinNordlund:ARFormation} demonstrated the formation of a small active region in a setup which did not impose initially a coherent flux bundle emerging though the bottom boundary (located at $z=-20$ Mm). Instead they imposed horizontal flux in inflow regions and the formation of a bipolar group occurred as consequence of convective scale selection within the simulation domain. On the one hand, this indicates a substantial degree of robustness with rather weak dependence on the initial field configuration. In light of this, perhaps the photospheric field structure may not tell us much about subsurface field properties. On the other hand, the choice of imposing a uniform horizontal field (with constant strength and orientation over the course of two days) in upflows crossing the bottom boundary implicitly assumes the existence of a large-scale, coherent magnetic structure at depths below $20$ Mm.  Recently \\citet{Brandenburg:NEMPISpots} proposed a mechanism for the near surface field amplification of pre-existing flux based on the so called negative magnetic pressure instability (NEMPI) that leads to super-equipartition field strengths regardless of subsurface conditions. While it is difficult to quantify if this mechanism also contributes to amplification in our simulations, the asymmetry of observed spots on the sun and the empirical fact the spot formation happens during or just after flux emergence suggests that NEMPI is not the sole mechanism for spot formation.  A more detailed comparison with flow observations during an active region formation phase is needed to put better constraints on the subsurface magnetic field and flow structure responsible for flux emergence. The numerical model presented here is a prototype simulation intended to highlight how a combination of upflows and field-aligned flows can influence the flux emergence and subsequent sunspot formation process. A detailed study comparing flows in the simulation presented here (and other comparable simulations) with flows inferred through helioseismology is work in progress."
    },
    "1402/1402.0316_arXiv.txt": {
        "abstract": "We have performed a 2.5 dimensional magnetohydrodynamic simulation that resolves the propagation and dissipation of Alfv\\'{e}n waves in the solar atmosphere. Alfv\\'{e}nic fluctuations are introduced on the bottom boundary of the extremely large simulation box that ranges from the photosphere to far above the solar wind acceleration region. Our model is ab initio in the sense that no corona and no wind are assumed initially. The numerical experiment reveals the quasi-steady solution that has the transition from the cool to the hot atmosphere and the emergence of the high speed wind. The global structure of the resulting hot wind solution fairly well agree with the coronal and the solar wind structure inferred from observations. The purpose of this study is to complement the previous paper by \\cite{2012ApJ...749....8M} and describe the more detailed results and the analysis method. These results include the dynamics of the transition region and the more precisely measured heating rate in the atmosphere. Particularly, the spatial distribution of the heating rate helps us to interpret the actual heating mechanisms in the numerical simulation. Our estimation method of heating rate turned out to be a good measure for dissipation of Alfv\\'{e}n waves and low beta fast waves. ",
        "introduction": "For years, extensive studies have been carried out on the coronal heating problem. The hot corona lies above the cool photosphere \\citep{1943ZA.....22...30E}, which can not be explained by thermal processes. Instead the mechanical energy injection from the surface convection motion is expected to maintain the hot corona by means of magnetic field. The coupling between the convection and the magnetic field transports free magnetic energy that will be converted to thermal energy in the upper atmosphere. To resolve the coronal heating problem, we should specify the physics that are responsible for releasing the free magnetic energy.  In this paper, we will focus on the atmosphere above the solar pole where the fast solar wind emanates. The magnetic field lines above the solar pole open up to the interplanetary space. One of the most plausible carriers of the magnetic free energy in the open field region is Alfv\\'{e}n waves, which was first proposed by \\cite{1947MNRAS.107..211A}. Recent observations have been accumulating several pieces of evidence of the existence of the propagating Alfv\\'{e}n waves in the solar atmosphere \\citep{2011ApJ...736L..24O}. Propagation and dissipation of Alfv\\'{e}n waves in the solar atmosphere is quite complicated since the gravity and the magnetic field make highly inhomogeneous atmospheres. For 1 dimensional configuration, nonlinear steepening of Alfv\\'{e}n waves and the subsequent shock formation play an important role in wave dissipation \\citep{1982SoPh...75...35H,1999ApJ...514..493K,2010ApJ...710.1857M}. Plenty of wave modes will emerge when we consider two dimensional configuration \\citep{1983PASJ...35..263S,2002ApJ...564..508R,2003ApJ...599..626B}.   All the above numerical simulations adopt simplified energy equation. In the coronal loop, mechanical heating is balanced by the radiative cooling and thermal conduction. Enthalpy flux due to the existence of the solar wind is also important when we consider the open field structure \\citep{1986JGR....91.4111H}. Therefore in order to obtain the atmospheric structure above the solar pole, we should treat the corona and the solar wind in a self-consistent way. One of the most important aspects of the self-consistent models is that these models can determine the mass loss rate from the sun given the boundary condition at the surface. If we can construct the appropriate mass loss model from the sun, the model could be very useful even for the stellar or the planetary atmospheres.  Pioneering works by \\cite{1982ApJ...259..779H,1982ApJ...259..767H} treated the corona and the solar wind simultaneously in a self-consistent way. Instead of using the artificial heating function, \\cite{1986JGR....91.4111H,2007ApJS..171..520C} adopted the physically based heating model from phenomenological turbulent theories. Heating models based on acoustic shocks \\citep{2002ApJ...578..598S} and MHD shocks \\citep{2004MNRAS.349.1227S} are also investigated. Time steady condition is relaxed to demonstrate the self-consistent reproduction of the hot coronal wind with fully dynamical 1D \\citep{2005ApJ...632L..49S,2006JGRA..11106101S} and 2D \\citep{2012ApJ...749....8M} (hereafter referred to as MS12) MHD simulations.  MS12 suggested that the shock heating is dominant heating mechanisms in the chromosphere and the coronal bottom while the turbulent heating is important in the solar wind acceleration region. However, there are two problems in the way how MS12 estimate the heating rate in their simulation. First, MS12 used dimensional analysis based on turbulent theory when they estimate the incompressible heating rate. This method can include large uncertainty, which allows MS12 only to estimate the order of magnitude of the heating rate. Second, the dimensional analysis does not have ability to estimate the temporal and spatial variability of the heating rate. In order to overcome the problems in MS12, we have developed an alternative method to estimate the heating rate. The purpose of the present paper is to describe more detailed features on heating mechanisms by using the new estimation method. Moreover, we will describe detailed dynamic features which are not shown in MS12.   We begin in section \\ref{sec:model} by describing our numerical models and corresponding assumptions. The time-averaged quasi-steady solution is shown in section \\ref{sec:solution} while the dynamic features relating to the coupling between waves and the transition region shall be discussed in section \\ref{sec:dynamics}. Then the basic idea to derive the heating rate as a function of space and time is given in section \\ref{sec:data_analysis}, although the detailed method for discretization and test simulations are explained in appendix \\ref{app:method}. In section \\ref{sec:heating}, we shall show the total amount and the spatial distribution of heating rate and discuss the possible mechanisms that actually happened in the numerical simulation. We continue in section \\ref{sec:discussion} by discussing the validity of the heating mechanisms found in our numerical simulation, while in section \\ref{sec:summary} we summarize.   \\section[]{Model \\& Assumptions}\\label{sec:model}  All the numerical methods to integrate the basic equations in this paper are the same as that used in MS12. We will describe the method in more detailed way here. Throughout the paper, we will assume single fluid MHD to describe a flux tube extended from the photosphere to the interplanetary space. Although the photosphere and the chromosphere are partially ionized, the use of single fluid approximation can be justified by frequent collisions between protons and neutrals \\citep[e.g.,][]{2008MNRAS.385.2269P}. Above the upper corona ($r-r_{\\rm s} > r_{\\rm s}$ or 7.0$\\times$10$^2$ Mm), proton and electron will decouple due to weak collisionality. Even in that case, we continue to use single fluid MHD equations for simplicity. Since we do not include any explicit dissipation terms in our equations, all the dissipations come from shocks and discretization errors. The method to estimate the heating rate shall be discussed in section \\ref{sec:data_analysis}. In addition to ideal MHD formulation, we include the gravity from the Sun. The solar gravity produces a highly stratified atmosphere , which inevitably leads initially linear waves to nonlinear waves. We consider a flux tube at $\\theta=\\pi/2$ in the spherical coordinate. We rotate the axis of the spherical coordinate so that the pole of the spherical coordinate lies on the equator of the sun. This helps us to avoid the numerical difficulty in the polar region of the spherical coordinate when we treat polar region of the sun. We neglect rotation of the sun as well as macro-scale magnetic field in this study. We impose the translational symmetry in the $\\theta$ direction ($\\partial _\\theta = 0$) so that all the variables depend on radius ($r$), azimuthal angle ($\\phi$), and time ($t$). Perturbations in $\\theta$ direction is purely incompressible, namely Alfv\\'{e}n mode, and perturbations of $r$ \\& $\\phi$ components consist of fast and slow MHD modes (see section \\ref{sec:dynamics}). Then our basic equations can be described as follows.  \\begin{eqnarray} {\\partial {\\cal U} \\over \\partial t} + {\\partial {\\cal F} \\over \\partial r} + {\\partial {\\cal G} \\over \\partial \\phi} = {\\cal S}, \\end{eqnarray}  \\begin{eqnarray} {\\cal U} = \\left( \\begin{array}{c} r^2 \\rho \\\\ r^2 \\rho V_r \\\\ r^2 \\rho V_\\theta \\\\ r^2 \\rho V_\\phi \\\\ r^2 B_r \\\\ rB_\\theta \\\\ rB_\\phi \\\\ r^2 E \\end{array}     \\right) \\label{eq:cnsvar} \\equiv \\left( \\begin{array}{c} {\\cal R}\\\\ {\\cal M}_r\\\\ {\\cal M}_\\theta\\\\ {\\cal M}_\\phi\\\\ {\\cal B}_r\\\\ {\\cal B}_\\theta\\\\ {\\cal B}_\\phi\\\\ {\\cal E} \\end{array}     \\right) , \\end{eqnarray}  \\begin{eqnarray} {\\cal S} = \\left( \\begin{array}{c} 0 \\\\ \\left[ - B_\\perp^2 + \\rho V_\\perp^2 + 2 P_{\\rm T} - {\\rho GM/r} \\right] \\times r \\\\ \\left( B_{r} B_\\theta - \\rho V_{r} V_\\theta \\right) \\times r\\\\ \\left( B_{r} B_\\phi - \\rho V_{r} V_\\phi \\right) \\times r\\\\ 0 \\\\ 0 \\\\ 0 \\\\ -\\rho V_{r} G M + Q_{\\rm cnd} + Q_{\\rm rad} \\end{array}     \\right) , \\end{eqnarray}  \\begin{eqnarray} {\\cal F} = \\left( \\begin{array}{c} \\rho V_{r} \\times r^2 \\\\ \\left( \\rho V_{r}^2 + P_{\\rm T} - B_{r}^2\\right) \\times r^2 \\\\ \\left( \\rho V_\\theta V_{r}  - B_{\\theta} B_{r} \\right) \\times r^2 \\\\ \\left( \\rho V_\\phi V_{r}  - B_{\\phi} B_{r} \\right) \\times r^2 \\\\ 0 \\\\ \\left( V_{r} B_\\theta - V_\\theta B_{r} \\right) \\times r \\\\ \\left( V_{r} B_\\phi - V_\\phi B_{r} \\right) \\times r \\\\ \\left[ \\left( E + P_{\\rm T} \\right) V_{r} - \\left( \\bmath{V} \\cdot \\bmath{B} \\right) B_{r} \\right] \\times r^2 \\end{array}     \\right) , \\end{eqnarray} \\begin{eqnarray} {\\cal G} = \\left( \\begin{array}{c} \\rho V_\\phi \\times r \\\\ \\left( \\rho V_{r} V_\\phi - B_{r} B_\\phi \\right) \\times r \\\\ \\left( \\rho V_\\theta V_\\phi - B_\\theta B_\\phi \\right) \\times r \\\\ \\left( \\rho V_\\phi^2 + P_{\\rm T} - B_\\phi^2\\right) \\times r \\\\ \\left( V_\\phi B_{r} - V_{r} B_\\phi \\right) \\times r \\\\ V_\\phi B_\\theta - V_\\theta B_\\phi \\\\ 0 \\\\ \\left[ \\left( E + P_{\\rm T} \\right) V_\\phi - \\left( \\bmath{V} \\cdot \\bmath{B} \\right) B_\\phi \\right] \\times r \\end{array}     \\right) , \\end{eqnarray}  \\begin{eqnarray} P_{\\rm g} &=& {\\rho k_{\\rm B} T \\over \\mu m_{\\rm H}},\\\\ E &=& {1\\over2}\\rho V^2 +{P_{\\rm g}\\over (\\gamma-1)} + {B^2\\over 2}, \\\\ P_{\\rm T} &=& P_{\\rm g} + {B^2 \\over 2},\\\\ V_{\\perp}^2 &=& V_\\theta^2 + V_\\phi^2, \\\\ B_{\\perp}^2 &=& B_\\theta^2 + B_\\phi^2, \\end{eqnarray} where all the symbols have standard meanings except for magnetic field, $B$, which is normalized by $\\sqrt{4 \\pi}$. We adopt the specific heat ratio $\\gamma$ of $5/3$ for mono-atomic molecules. At around 6,000 K or less, almost all of hydrogens are neutrals while they are completely ionized at the temperature above 10$^4$ K. To mimic the real equation of states there, we define the mean molecular weight as a function of temperature to linearly connect fully ionized and neutral gases, \\begin{eqnarray} \\mu = \\left\\{ \\begin{array}{c} 1 ~~\\left( T < 6,000 ~{\\rm K} \\right) \\\\ \\\\ \\left[ 3.5 \\left( 1-4285.7/T \\right) \\right]^{-1} \\\\ \\left( 6,000 ~{\\rm K} < T < 10^4 ~{\\rm K} \\right) \\\\ \\\\ 0.5 ~~\\left( {\\rm otherwise} \\right) \\end{array} \\right. . \\end{eqnarray} As for the thermal conduction, we use the classical formulation for collisional plasma, \\begin{equation} Q_{\\rm cnd} = \\nabla \\cdot \\left( \\kappa_0 T^{5/2} { \\nabla T \\cdot \\bmath{B} \\over B^2} \\bmath{B} \\right). \\label{qcnd} \\end{equation} We choose $\\kappa_0 = 10^{-6}$ in cgs unit, reasonable value of electrons for fully ionized plasma in thermal equilibrium \\citep{1965RvPP....1..205B}. Heat transport across magnetic field lines is inhibited in equation (\\ref{qcnd}). In the present study, we use this classical limit throughout all space, although the deviation from the the classical heat conduction could affect the temperature structure in the outer corona ($>10$ $r_{\\rm s}$ or 7.0 $\\times$ 10$^3$ Mm) \\citep[see][]{1978RvGSP..16..689H}. We also introduce radiative cooling as follows. \\begin{eqnarray} Q_{\\rm rad} = \\left\\{ \\begin{array}{lc} - n^2 \\Lambda \\left( T \\right) & ~(T>T_{\\rm c} ~\\&~ \\rho< \\rho _{\\rm c})\\\\ - n n_{\\rm c} \\Lambda \\left( T_{\\rm c} \\right) &({\\rm otherwise}) \\label{eq:radloss} \\end{array} \\right. , \\label{qrad} \\end{eqnarray} \\begin{eqnarray} \\begin{array}{cllc} \\log _{10} \\Lambda \\left( T \\right) &=-&127.3 & \\\\ &+& 56.57 & \\log _{10} T  \\\\ &-& 9.84 & (\\log _{10} T )^2 \\\\ &+& 0.5548 & (\\log _{10} T )^3 \\label{eq:lambda_func}, \\end{array} \\end{eqnarray} where $n=\\rho/m_{\\rm H}$, $T_{\\rm c} = 4 \\times 10^4$ K, and $\\rho_{\\rm c} = 4.9 \\times 10^{-17}$ g cm$^{-3}$. The term in the first line in equation (\\ref{eq:radloss}) represents the radiative cooling for optically thin plasma of the coronal abundance \\citep[e.g.,][]{1990A&AS...82..229L} fitted by polynomials. The radiative loss function is shown in figure \\ref{fig:lambda_func}. The term in the second line of equation (\\ref{eq:radloss}) is in proportion to $n$ and mimics the radiative cooling for optically thick plasma that was empirically derived by \\cite{1989ApJ...346.1010A}. Basically, the treatment of the thermal conduction and the radiative cooling in our study is the same one as that used in \\cite{2004ApJ...601L.107M,2005ApJ...632L..49S,2006JGRA..11106101S,2008ApJ...688..669A, 2010ApJ...712..494A}.  \\begin{figure}% \\begin{center} \\includegraphics[scale=1.0]{fig01.eps} \\caption{ The radiative loss function as a function of temperature. } \\label{fig:lambda_func} \\end{center} \\end{figure}   We set up our initial atmosphere by solving the equations of hydrostatic equilibrium without magnetic field. The initial atmospheres have 10$^4$ K all over the numerical domain, starting from the initial photospheric density of 10$^{-7}$ g cm$^{-3}$ at $r=r_{\\rm s}$. Above the height of 11 Mm ($\\sim$ 1.6 $\\times10^{-2}~r_{\\rm s}$), the density profiles are forced to be proportional to $r^{-2}$. This is simply because the exponential density drop causes numerical difficulties in the higher atmosphere. At the photospheric surface, we impose the locally concentrated magnetic field that would correspond to the kilo Gauss patch at the photosphere. The exact profile along $\\phi$ direction at the photosphere are shown with the solid line in Figure \\ref{initial_field}d. This initial profile is produced by the superposition of narrow positive Gaussian and wide negative Gaussian. The presence of the kilo Gauss patches is clearly confirmed by {\\it HINODE}/SOT even in the polar coronal hole where the high speed solar wind originates from \\citep{2008ApJ...688.1374T,2012ApJ...753..157S}. The magnetic field above the photosphere is extrapolated by using potential field approximation (Figure \\ref{initial_field}b). The super-radial expansion of the flux tube reduces the field strength from $\\sim$ 3$\\times$10$^3$ to 3 Gauss in the region below 4 Mm (Figure \\ref{initial_field}cd). Above $\\approx$ 4 Mm (6$\\times$10$^{-3}$ $r_{\\rm s}$), the field strength decreases in proportion to $r^{-2}$. Since all the open field lines start from a single cell in the current resolution, we have performed test simulations to check the resolution dependence in appendix \\ref{sec:res_test}. We found that as far as the energy flux at the corona is concerned, the results would not change drastically. The field strength 3 Gauss well below the TR could be weak compared with the usual polar field strength of $\\sim$ 10 Gauss. This is just a computational issue. Since we solve thermal conduction implicitly, the time step is usually determined by the Alfv\\'{e}n crossing time of the radial grid around the transition region ($\\sim 0.1$ sec in the current resolution). If we reduce the field strength, we can use a longer time step restricted by the CFL condition. Larger field strength will produce larger Poynting flux above, which can be potential source to heat the corona.  \\begin{figure}% \\begin{center} \\includegraphics[scale=1.0]{fig02.eps} \\caption{Magnetic field structure in our numerical simulation. (a) Mean magnetic field structure in the quasi-steady state below 10 Mm (1.4$\\times$10$^{-2}$ $r_{\\rm s}$). The black(gray) solid lines show open(closed) magnetic field lines. The black dashed lines indicate iso contours of plasma beta.The magnetic field lines (b) The initial magnetic field lines below 10 Mm (1.4$\\times$10$^{-2}$ $r_{\\rm s}$). The black(gray) solid lines represent the open(closed) field lines. (c) The radial magnetic field strength along $\\phi = 0$ as a function of height. (d) The radial magnetic field strength along the azimuthal direction. The solid, dashed, and dotted lines represent the field strength at $r-r_{\\rm s}=0,2,6$ Mm (0, 2.9$\\times$10$^{-3}$, 8.6$\\times$10$^{-3}$ $r_{\\rm s}$), respectively. The black/gray color means positive/negative value. } \\label{initial_field} \\end{center} \\end{figure}  The horizontal length of our numerical domain is 3,000 km at the photosphere and it expands in proportion to $r$. 32 grid points are allocated to resolve 3,000 km so that the horizontal spatial resolution is nearly 100 km at the photosphere. Periodic boundary conditions are posed on the horizontal boundaries. As for the radial direction, we start our simulations from the photosphere ($r=r_{\\rm s}$ or 7$\\times$10$^2$ Mm) to the interplanetary space ($r=1,674~r_{\\rm s}$ $> 7$AU or 10$^6$ Mm). The radial spatial resolution is 25 km at the photosphere. We increase the grid size non-uniformly ( 0.7 $r_{\\rm s}$ or 490 Mm at 7 AU) as we go up higher to cover the whole numerical domain by using 16,384 grid points. The total number of the radial grid points are doubled from the previous simulation by MS12 to better resolve the wave propagation in the low corona. Open boundary in radial direction could be implemented by using characteristic method, although it is not straightforward to use characteristic method with thermal conduction. Instead, we take lengthy radial domain to avoid the numerical reflection from the top boundary and pose zero-derivative boundary condition there. At the final time of our simulation, the thermal conduction front reaches 4 AU (8.6$\\times$10$^2$ r$_{\\rm s}$ or 6$\\times$10$^5$ Mm) while the solar wind reaches 1 AU (2.2$\\times$10$^2$ r$_{\\rm s}$ or 1.5$\\times$10$^5$ Mm). Therefore, any physical information propagating from the inner boundary can not reach the top boundary within the duration we considered. In this paper we focus on the wind structure in $r<20$ $r_{\\rm s}$ ($\\sim 0.1$ AU, 1.4$\\times$10$^4$ Mm), which is in the quasi-steady state after sufficient Alfv\\'{e}n crossing time.  At the inner boundary, all the variables are fixed to the initial value except for the velocity in $\\theta$ direction ($V_\\theta$). The spectra of $V_\\theta$ at the bottom boundary are prescribed to excite Alfv\\'{e}n waves. Throughout this paper, we will restrict our self to investigate the Alfv\\'{e}n mode (fluctuations in $\\theta$ direction). The one-sided power spectrum of $V_\\theta$ are defined by \\begin{eqnarray} P_\\nu = 2 \\lim_{{\\tau} \\to \\infty } {\\langle \\hat{V_\\theta}(\\nu)\\hat{V_\\theta}^*(\\nu) \\rangle \\over {\\tau} }, \\end{eqnarray} where $\\langle \\rangle$ here means temporal average over period $\\tau$, $\\nu$ is frequency, and $\\hat{V_\\theta}(\\nu)$ are complex Fourier components of $V_\\theta (t)$, \\begin{eqnarray} \\hat{V_\\theta}(\\nu) &=& \\int^{\\infty}_{-\\infty} V_\\theta (t) e^{-2\\pi i \\nu t} dt. \\end{eqnarray} Using the one-sided power spectrum, the total power of $V_\\theta$ can be described by \\begin{eqnarray} \\langle V_\\theta (t)^2 \\rangle &=& \\int^{\\infty}_{0} P_\\nu d\\nu . \\end{eqnarray} We will concentrate on the white noise case ($P_\\nu \\propto \\nu^0$) with the total power of (2.2 km s$^{-1}$)$^2$. This total power seems twice larger than the observed photospheric velocity \\citep{2010ApJ...716L..19M}. Since we fixed $B_{\\theta}$ to be zero at the bottom boundary, out going Els\\\"{a}sser variables, \\begin{eqnarray} z^-_{\\theta}\\equiv {1\\over 2} \\left( V_\\theta - {B_\\theta \\over \\sqrt{\\rho}} \\right), \\end{eqnarray} becomes the half of $V_\\theta$ specified above. Therefore, we adopt the total power of (2.2 km s$^{-1}$)$^2$ so that $\\langle(z^-_\\theta)^2 \\rangle=$(1.1 km s$^{-1}$)$^2$. The wave energy flux with this boundary condition ($\\rho\\sim10^{-7}$ g cm$^{-3}$, $z^-_{\\theta}\\sim1$ km s$^{-1}$, and V$_{\\rm A} \\sim 10$ km s$^{-1}$) becomes $\\sim$10$^9$ erg cm$^{-2}$ s$^{-1}$ at the photosphere, which ensures adequate energy supply provided dissipation works. The frequency of the fluctuations is restricted to range from 2.5 $\\times$ 10$^{-4}$ Hz (4,000 sec) to 2 $\\times$ 10$^{-2}$ Hz (50 sec) in this study. The waves are specified over the entire lower boundary.  Our numerical code is basically based on HLLD scheme, an approximate Riemann solver that have robustness and inexpensive numerical cost \\citep{2005JCoPh.208..315M}. We combined HLLD scheme and flux-CT method \\citep{2000JCoPh.161..605T,2005JCoPh.205..509G} to preserve the initial $\\nabla \\cdot \\bmath{B}$ within the rounding error. The initial $\\nabla \\cdot \\bmath{B}$ can be as small as the rounding error if we use the vector potential rather than the scalar potential to extrapolate the initial magnetic field. In order to avoid the negative pressure, we modified the energy equation in the similar way as \\cite{1999JCoPh.149..270B}. This modification violates the energy conservation slightly at the level of discretization errors. Using the MUSCL interpolation and the 2nd order Runge-Kutta integration methods, our numerical code achieves the 2nd order accuracy in both space and time. ",
        "conclusions": "\\label{sec:discussion}  Thanks to the new analysis method described in section \\ref{sec:data_analysis} and appendix \\ref{app:method}, we can derive more accurate profile of the heating rate than that was reported by MS12. Although the numerical simulation shown here is almost the same as that of MS12, the new analysis method and higher numerical resolution lead us two different interpretations about heating mechanisms in the corona.  First, MS12 suggested that the shock dissipation and MHD turbulence are the main heating mechanisms at the coronal bottom. We concluded that, instead of MHD turbulence, Alfv\\'{e}n waves with curved wedge shape originate from the chromospheric fast shocks is main heating mechanisms at the coronal bottom. Phase mixing due to the concave transition region may also stimulate the coronal heating. The heating rate derived in this paper is more reliable than that of MS12. Next, MS12 suggested that MHD turbulence is the dominant heating mechanisms in the solar wind. The finer radial resolution reveals that we do not achieve the shallower power spectrum like Kolmogorov-type turbulence. We confirm that the magnetic energy are converted to the thermal energy through incompressible processes, although it is difficult to distinguish what kind of heating mechanisms operate in the solar wind.  We derive heating rate in our numerical simulation that inevitably has numerical dissipation larger than that of the realistic solar atmosphere. Therefore we should discuss the validity or the robustness of the heating mechanisms found in our simulations, although we do need further numerical experiments to investigate whether our results depend on numerical resolution. Shock heating rate would not depend on the grid size and microscopic physics like viscosity or resistivity as far as Rankine-Hugoniot relations are satisfied. Since our numerical scheme (HLLD scheme) is conservative and meditated to capture MHD shocks, the shock heating in our numerical simulation should not be just a numerical heating but a physical one.  Heating events at the coronal bottom are associated with curved wedge shape structure described in section \\ref{sec:heating}. This structure would be Alfv\\'{e}n wave front with large current density that was stored in the chromospheric fast shocks. Whether this structure keeps large current density or not might depend on the numerical resolution. The dynamic simulations with high resolution are needed to justify this heating mechanism.  The heating by phase mixing depends on viscosity or resistivity so that our results could be affected by numerical resolution. The phase mixing in our simulation turns on because of the horizontal phase speed difference that is created by the collision between the transition region and shock waves. Then the duration that phase mixing operates is finite and determined by the Alfv\\'{e}n crossing time between peaks and troughs of the transition region. The finite duration for phase mixing results in the finite thickness of the current density. Therefore this will prevent phase mixing from cascading to the dissipation scale that is determined by microscopic physics, although, in our simulation, Alfv\\'{e}n waves dissipate due to insufficient grid numbers in $\\phi$ direction. Therefore there is a possibility that the phase mixing discussed here may be weaken to operate if we increase grid numbers. Extension to 3 dimensional simulation could stimulate phase mixing by increasing the effective viscosity or resistivity through instability due to velocity or magnetic field shear. Or inhomogeneity in $\\theta$ direction makes phase mixing less effective as was suggested by \\cite{1991ApJ...376..355P}.  In the present study, we only consider shear Alfv\\'{e}n fluctuations ($V_{\\theta}$) as the energy injection from the photosphere. However the realistic fluctuations would also be associated with kink and torsional types in the complex solar atmosphere without translational symmetry ($\\partial _\\theta = 0$). Torsional Alfv\\'{e}n waves could be generated by the vortex flows observed in the photosphere \\citep{1988Natur.335..238B}, and the connection between the vortex flows and swirling motions in the upper atmosphere are observed recently \\citep{2012Natur.486..505W}. Since the difference in nonlinear behaviors between torsional and shear mode is discussed by \\cite{2011A&A...526A..80V}, dynamic behaviors could be affected if we include torsional Alfv\\'{e}n waves. This aspect will be explored in future work by using axisymetric 2.5 dimensional simulations or full 3 dimensional simulations.  Since we only prescribe Alfv\\'{e}nic fluctuations that cannot be linearly converted into compressible mode, the nonlinear development of Alfv\\'{e}n waves is important to heat the atmosphere. The wave nonlinearity $V_{\\theta}/V_{\\rm A}$ is proportional to $\\rho^{1/4}B^{-1}$ if we consider the wave energy flux conservation in an static atmosphere, $\\rho V_{\\theta}^2 V_{\\rm A} A=$ const, where $A$ is the cross section of the flux tube. This means that not only the density structure but also the expansion rate of the magnetic field are important for the development of the wave nonlinearity. The previous 1 dimensional simulations artificially defined the cross section of flux tubes. The magnetic field in our model has nearly the potential field structure similar to that in \\cite{2003ApJ...599..626B}, which makes our simulation more realistic than the 1 dimensional simulations. Our model, however, cannot describe the non-magnetic atmosphere below the magnetic canopy, since potential field structure needs the surrounding magnetic pressure to support the strong magnetic flux tube. For more realistic simulation, we should implement the magnetohydrostatic atmosphere in the similar way \\cite{2005ApJ...631.1270H} did to investigate the compressible wave propagation.  From the chromospheric observations with HINODE/SOT, \\cite{2007PASJ...59S.655D} suggested that there are two types of spicules: Type I spicules exhibit up-and-down motion and have relatively slow ascending speed while Type II spicules fade out and often have high speed upward motion larger than 100 km s$^{-1}$. If we regard the up-and-down motion of the transition region as spicules, the radial velocity of our spicules is usually 60 km s$^{-1}$ at most. Therefore the collision between shock waves and the transition region could not lift up the plasma so rapidly as Type II spicules do and other mechanisms will be needed \\citep[e.g.][]{2011ApJ...736....9M}. It should be noted however that the existence of the Type II spicules is still under debate so far \\citep{2012ApJ...750...16Z,2012ApJ...759...18P}.     We have implemented a 2.5 dimensional MHD simulation that resolves the propagation and the dissipation of the Alfv\\'{e}n waves in the solar atmosphere. The hot corona and the high speed solar wind are reproduced as a natural consequence of the Alfv\\'{e}n wave injection from the photosphere, as was shown in MS12. However, the detailed analysis of the heating rate leads us to the different interpletation as follows. \\begin{itemize} \\item Alfv\\'{e}n waves with curved wedge shape generated by the chromospheric fast shocks heat the coronal bottom while MS12 suggested heating by MHD turbulence. \\item Anisotropic turbulent cascade that was found in the simulation of MS12 turned out to disappear when increasing radial resolution. The heating in the solar wind operates through the magnetic energy conversion with incompressible process, though we could not distinguish the specific heating mechanisms so far. \\end{itemize}"
    },
    "1402/1402.6126_arXiv.txt": {
        "abstract": "RS~Oph is a recurrent symbiotic nova that undergoes nova-like outbursts on a time scale of 20 years. Its two last eruptions (1985 and 2006) were subject of intensive multiwavelengths observational campaign from the X-rays to the radio. This contribution aims to determine physical parameters and the ionization structure of the nova from its explosion to the first emergence of the supersoft X-rays (day 26) by using the method of multiwavelength modelling the SED. From the very beginning of the eruption, the model SED revealed the presence of both a strong stellar and nebular component of radiation in the spectrum. During the first 4 days, the nova evinced a biconical ionization structure. The $\\sim 8200$\\,K warm and 160--200\\ro\\ extended pseudophotosphere encompassed the white dwarf (WD) around its equator to the latitude $>40^{\\circ}$. The remaining space around the WD's poles was ionized, producing a strong nebular continuum with the emission measure $\\textsl{EM} \\sim 2.3 \\times 10^{62}$\\cmt\\ via the fast wind from the WD. The luminosity of the burning WD was highly super-Eddington for the whole investigated period. The wind mass loss at rates of $10^{-4} - 10^{-5}$\\myr\\ and the presence of jets suggest an accretion throughout a disk at a high rate, which can help to sustain the super-Eddington luminosity of the accretor for a long time. ",
        "introduction": "The star RS~Ophiuchi (RS~Oph) is located in the direction nearby to the galactic center ($l_{\\rm II} = +19^{\\circ}.8 ~ b_{\\rm II} = +10^{\\circ}.4$) at a distance $d = 1.6 \\pm 0.3\\kpc$ \\citep[][]{hjell+86,bode87,barry08a}. The light of RS~Oph is significantly attenuated by interstellar extinction. \\cite{hjell+86} determined from \\hi\\ 21\\,cm measurements the interstellar absorbing column, $N_{\\rm H} = (2.4\\pm 0.3)\\times 10^{21}$\\,\\cmd, and \\cite{snijders87} derived the color excess $E_{\\rm B-V} = 0.73\\pm 0.1$\\,mag from ultraviolet \\textsl{IUE} spectra. According to the relationship between $N_{\\rm H}$ and $E_{\\rm B-V}$ \\citep[e.g.][]{d+s94}, both these parameters are consistent.  RS~Oph is a recurrent symbiotic nova comprising a late-type, K7\\,III, giant \\citep{ms99} and a WD with a mass close to the Chandrasekhar limit \\citep[e.g.][and references therein]{bode87}, in a 456-day orbit \\citep{f+00,brandi+09}. Its nova-like outbursts are characterized with brightening by about 7\\,mag and a recurrence period of about 20 years. Historically, 6 eruptions have been recorded unambiguously. The first one in 1898 and the last one on 2006 February 12.83 \\citep[][]{evans88,narumi+06}. The recurrence period of approximately 20 years and a bright peak magnitude, $V = 4-5$, made RS~Oph a good target for multifrequency observational campaigns \\citep[e.g.][and references therein]{evans07}.  At the beginning of the last two outbursts of 1985 and 2006, observations revealed a non-spherical shape of the nova ejecta, (i) on the radio maps \\citep[e.g.][]{taylor+89,tob+06,rupen+08,sok+08}, (ii) by interferometric technique in the near and mid-IR \\citep[e.g.][]{monnier+06,lane+07,chesneau+07,barry08b}, (iii) by the \\textsl{HST} imaging in the optical \\citep{bode+07}, (iv) by the \\textsl{Chandra} X-ray Observatory detected as an extended X-ray emission elongated in the line with the extended infrared emission \\citep[][]{luna+09} and (v) by the asymmetric line profiles measured on day 13.9, interpreted by \\cite{drake+09} as a result of the shock collimation due to equatorial circumstellar density enhancement. Particularly, highly collimated jet-like ejection was indicated by the optical spectroscopy during the first 30 days after the eruption \\citep[][]{sk+08a,banerjee}.  The 2006 outburst was intensively monitored with the {\\em Swift} and the {\\em Rossi X-ray Timing Explorer}. Prior to day 26, the behaviour of the X-ray flux was described by the evolution of shock systems established as the high-velocity ejecta impacted the red giant wind \\citep[][]{bode+06,sok+06}. The evolving shock in the 2006 outburst of RS~Oph was analyzed in detail by \\cite{ness+09} using the high-resolution X-ray grating observations with \\textsl{Chandra} and \\textsl{XMM-Newton}. The X-ray observations demonstrated that the supersoft X-ray phase of nova RS~Oph started rapidly from day 26 \\citep[see also][]{osborne+11}.  The detailed spectral evolution of the 2006 outburst in the optical was reported by \\cite{iijima09}. He suggested that the profiles of the prominent emission lines resulted from a rapidly expanding ($\\sim 1000$\\kms), low density part and a slowly moving ($\\sim 100$\\kms) high density part in the ejecta, which are probably related to the expansion along the polar and the equatorial regions, respectively.  In this contribution I aim to determine physical parameters of the nova RS~Oph from its explosion to the first emergence of the supersoft X-rays (day 26) by using the method of multiwavelength modelling the SED. Based on the results from spectral fits, I also explore the ionization structure of the nova during the first 4 days of its outburst. Section~2 introduces multi-band observations at selected days, and the results of their SED-fitting analysis are given in Sect.~3. Discussion and summary are found in Sects.~4 and 5, respectively. ",
        "conclusions": "\\subsection{Biconical ionization structure during the first 4 days}  Multiwavelength modelling of the SED revealed the simultaneous presence of a warm stellar ($F_{\\rm h}(\\lambda)$) and nebular ($F_{\\rm n}(\\lambda)$) components of radiation in the spectrum of RS~Oph during the first four days after its explosion. The stellar component is characterized with $T_{\\rm h} = 9600-6700$\\,K and $L_{\\rm h} = 7.2 - 1.9 \\times 10^{38}$\\es, which corresponds to production of $4.4\\times 10^{45} - 1.1\\times 10^{42}$ photons per second that are capable of ionizing hydrogen. However, the nebular component is characterized with $\\textsl{EM} \\sim 2.3 \\times 10^{62}$\\cmt, which corresponds to $\\sim 3 \\times 10^{49}$ recombinations per second $(\\alpha_{\\rm B}({\\rm H},22000\\,K) = 1.33\\times 10^{-13})$, and thus cannot be generated by the observed stellar source of radiation ($L_{\\rm ph} \\lll \\alpha_{\\rm B}\\textsl{EM}$; see Eq.~(\\ref{eq:lph})). This implies the presence of a hot ionizing source in the system, which is not seen directly by the outer observer. Such properties of radiative components were revealed in the spectrum of all symbiotic stars with a high orbital inclination during their outbursts. The corresponding spectrum was called a two-temperature-type \\citep[see Sect.~5.3.4 of][]{sk05}. This puts constraints for the ionization and geometrical structure of the circumstellar matter surrounding the burning WD. According to \\cite{sk05} the source of the warm stellar component can be ascribed to the WD's pseudophotosphere that is represented by the outer flared rim of an optically thick disk-like formation around the hot star. The circumstellar material located within the remainder of the sphere is ionized, and thus converts a fraction of the hot star (here the burning WD) radiation into the strong nebular emission (see Fig.~\\ref{fig:sketch}).  \\subsubsection{Nature of the ionized region}  Ever since the very beginning of the nova explosion (day 1.16), very broad emission lines of hydrogen were observed in the optical spectrum \\citep[e.g.][]{banerjee}. \\cite{sk+08b} fitted the broad wings of \\ha\\ from day 1.38 by the wind model of \\cite{sk06}, in which a fraction of the wind from the central star is blocked by a disk around the WD's equator, and thus causes its bipolarity. The very high level of the continuum around the \\ha\\ line of $2.5\\times 10^{-10}$\\ecsa\\ (see Fig.~\\ref{fig:seds}, left top panel) yields its luminosity, $L_{\\alpha}\\sim 4900$\\lo\\ and the mass loss rate of a few $\\times\\,10^{-4}$\\myr\\ \\citep[][]{sk+08b}. Consequently, the \\ha\\ luminosity was generated in a volume with the emission measure \\textsl{EM}$_{\\alpha} = L_{\\alpha}/\\varepsilon_{\\alpha} \\sim 10^{62}(d/1.6\\kpc)^2$\\cmt\\ for $\\varepsilon_{\\alpha}(20000\\,K) = 1.83\\times 10^{-25}$\\,erg\\,cm$^{3}$\\,s$^{-1}$, which is consistent with that of the nebular continuum. \\footnote{Its somewhat smaller value can be caused by a larger optical depth in the line and by contributions from singly and doubly ionized helium to the nebular continuum, which were not subtracted.} As both the nebular continuum and the hydrogen emission lines originate in the same region, the H\\I\\I\\ zone, producing the strong nebular component of radiation in the SED, can be ascribed to the fast ionized wind from the outbursting WD.  In addition, the biconical structure and the nature of the ionized zone was also supported by the presence of satellite emission components located at $\\pm (2400-2500)$\\kms\\ with respect to the main cores of hydrogen lines, measured from the very beginning of the outburst and being recognizable to day $\\sim 30$ \\citep[][]{sk+08b,banerjee}. These emission features can be interpreted as a result of highly collimated bipolar mass outflow. This was confirmed by the radio 43\\,GHz image from day 55 that showed directly highly collimated outflows produced by free-free emission from the thermal plasma \\citep[][]{sok+08}. The presence of bipolar jets requires the presence of the inner disk, throughout which sufficient amount of mass must be accreted to balance the outflow via the jets \\citep[e.g.][]{livio+03}. Thus the presence of a large disk-like formation encompassing the burning WD could also be responsible for keeping the super-Eddington luminosity of the burning WD for a long time. \\begin{figure*}[!t] \\centering \\begin{center} \\resizebox{11cm}{!}{\\includegraphics[angle=-90]{rs1-4.eps}} \\end{center} \\caption[]{ A sketch of the ionization structure of RS~Oph during the first 4 days after its explosion as seen on a cut perpendicular to the orbital plane containing the burning WD (the black circle). The neutral (H\\,\\I) zone is optically thick to the distance $R_{\\rm D}$ from the WD, which represents the radius of the warm pseudophotosphere (heavy solid line). The ionized (H\\,\\I\\I) zone has the opening angle $\\theta_{\\rm HII}$. Dimensions in the figure are given by the model SED and the elements of the spectroscopic orbit according to \\cite{brandi+09}. } \\label{fig:sketch} \\end{figure*}  \\subsubsection{Consequences of the bipolar structure}  (i) The bolometric luminosity of the burning WD cannot be determined unambiguously during this period of the nova evolution. Only its lower limit can be estimated from the components of radiation determining the observed two-temperature-type of the spectrum. The disk-like shell around the WD converts a fraction of its radiation into the warm stellar component (here approximated by blackbody). Its luminosity at day $\\sim 1$ was $7.3\\times 10^{38}$\\es\\ (Table~2). Another fraction of the burning WD's luminosity is reprocessed by ionization/recombination events into the nebular component. Its luminosity can be expressed as (see Eq.~(8) in Paper~I) \\begin{equation} L_{\\rm n} =  \\textsl{EM} \\int_{\\lambda > 912}^{\\infty} \\varepsilon_{\\lambda}(T_{\\rm e}) {\\rm d}\\lambda , \\label{eq:ln} \\end{equation} where the nebula is assumed to radiate under conditions of Case~B. The \\textsl{EM}(day $\\sim 1) = 3.2\\times 10^{62}$\\cmt\\ corresponds to $L_{\\rm n} = 2.6\\times 10^{38}$\\es. This also represents a lower limit, because a fraction of ionizing photons is not detectable (they are not converted into the nebular photons, i.e. $L_{\\rm ph} > \\alpha_{\\rm B}\\textsl{EM}$, see Appendix~A) and, in addition, the nebula can be partially optically thick (i.e. the measured \\textsl{EM} is smaller than that originally created by ionizations). Thus the bolometric luminosity of the burning WD is \\begin{equation} L_{\\rm h}^{\\rm bol} > L_{\\rm h} + L_{\\rm n} = 9.9\\times 10^{38}\\,{\\rm erg}\\,{\\rm s}^{-1} \\label{eq:lbol} \\end{equation} at day $\\sim 1$. Note that such adding of luminosities from different sources to the bolometric luminosity is possible only in the case of the two-temperature-type of the spectrum (i.e. the warm pseudophotosphere is far below the capability of giving rise to the observed nebular component of radiation).  (ii) According to the bipolar ionization structure of the burning WD, the vertical extension of the warm pseudophotosphere is limited by the orbital inclination to get hidden the ionizing source in direction of the observer. Its luminosity can be written as \\begin{equation} S(R_{\\rm D},\\theta_{\\rm D}) \\sigma T_{\\rm eff}^4 = 4\\pi (R_{\\rm h}^{\\rm eff})^2\\sigma T_{\\rm eff}^4 , \\label{eq:wlum} \\end{equation} where the surface $S(R_{\\rm D},\\theta_{\\rm D})$ of the pseudophotosphere is represented by a belt on a sphere with the radius $R_{\\rm D}$, located around its equator and extended to the latitude $\\pm \\theta_{\\rm D}/2$ (see Fig.~\\ref{fig:sketch}). $R_{\\rm h}^{\\rm eff}$ is its effective radius (see Eq.~(6) in Paper~I). Then the radial extension of such the pseudophotosphere can be expressed as \\begin{equation} R_{\\rm D} = \\frac{R_{\\rm h}^{\\rm eff}} {\\sqrt{\\sin(\\theta_{\\rm D}/2)}}. \\label{eq:rd} \\end{equation} For $R_{\\rm h}^{\\rm eff} \\sim 160$\\ro\\ (Table~2) and the limiting case $\\theta_{\\rm D}/2 = 90^{\\circ} - i$, where $i = 50^{\\circ}$, we get $R_{\\rm D} = 200$\\ro. Other possibilities, $40^{\\circ} < \\theta_{\\rm D}/2 < 90^{\\circ}$ yield $200 > R_{\\rm D} > 160$\\ro.  \\subsubsection{Comparison with other measurements}  At day 3.8 after the nova explosion, \\cite{barry08b} used the KIN to find the structure of the circumstellar material around the star. They determined the angular size of the mid-infrared continuum emitting material to 4.0--6.2\\,mas, which corresponds to linear size of 6.4--9.9(d/1.6\\kpc)\\,AU. As the N-band region is dominated by the nebular emission (see Fig.~\\ref{fig:seds}), which is due to the fast ionized wind from the WD (see Sect.~4.1.1), its maximum extension (13\\,AU for the velocity of 3000\\kms\\ and 3.8 days) is consistent with that detected by the KIN. The measurements obtained by the KIN in the inner spatial regime showed evidence of enhanced neutral atomic hydrogen emission and atomic metals located near the WD. This could reflect the H\\,\\I\\ zone in the ionization structure inferred from the model SED (Fig.~\\ref{fig:sketch}).  \\cite{monnier+06} measured the near-IR size of RS~Oph using the IOTA and Keck interferometer at the same day. They found a significantly asymmetric emission with the characteristic size of FWHM$\\sim 3$\\,mas. Observations for the following 60 days did not indicate any change, suggesting thus that the IR emission is due to the $b-f$ and $f-f$ transitions in thermal plasma rather than from a hot gas in the expanding shock as considered in current popular models. Also these results -- the size and the nature of the IR emission -- are consistent with the model SED here.  Using the AMBER/VLT interferometer, \\cite{chesneau+07} measured the extension of the milliarcsecond-scale emission in the K band continuum, in the Br${\\gamma}$ and He\\,\\I\\,2.06\\,$\\mu$m lines at day 5.5. Observations showed that the emission from all investigated spectral regions is highly flattened with the ratio $b/a\\sim 0.6$, having the same position angle and a typical Gaussian extension of $3.1\\times 1.9$\\,mas for the continuum.  Thus, it is possible to conclude that the biconical ionization structure of the circumstellar matter around the WD in RS~Oph, as inferred from multiwavelength modelling the SED, is supported by direct interferometric measurements using IOTA, Keck and AMBER/VLTI performed during a few days after the eruption.  \\subsection{Sudden transition of the SED maximum to the UV}  The model SED for day 7.3 showed that the shift of the SED maximum from the optical to the far-UV (or even beyond) happened during a relatively short time of $< 3.5$ days (see Fig.~\\ref{fig:seds}). A sudden decrease of the optical depth causing the shrinkage of the inflated WD from $\\sim 160$\\ro\\ to less than 2.1\\ro\\ (Table~2) would require removal of a significant amount of dense, optically thick, material in the line of sight. The corresponding time of such the change should be proportional to the dynamical time scale of the inflated WD, \\begin{equation} t_{\\rm dyn} = (G \\rho)^{-1/2} , \\label{eq:tdyn} \\end{equation} where $\\rho$ is the density of the matter giving rise to the gravitational force. So, for the WD mass of $1.3$\\mo, its effective radius 160\\ro\\ would collapse during $t_{\\rm dyn} = 67$ days, which is by a factor of $\\sim 20$ larger than the time scale of the observed change in the effective radius. Therefore, a gradual shrinking of the WD pseudophotosphere due to the expansion of the outflowing material cannot be the cause of its sudden disappearance between day 3.8 and 7.3.  The sudden change in the SED can be a result of a change in the position of the ionization boundary with respect to the line of sight. If the opening angle of the H\\,\\I\\ zone $\\theta_{\\rm D}$ (see Fig.~\\ref{fig:sketch}) shrinks so that $\\theta_{\\rm D}/2 < 90^{\\circ} - i$, the burning WD will rise up the horizon represented by the H\\,\\I/H\\,\\I\\I\\ boundary. A shrinkage/enlargement of the $\\theta_{\\rm D}/\\theta_{\\rm HII}$ angle can be a result of a decreasing mass loss rate from the WD. Recently, \\cite{cask12} found that the ionization structure of the hot components in symbiotic binaries during outbursts is similar to that sketched in the figure~\\ref{fig:sketch} here (see their Figs.~1 and 6). In their model, the opening angle of the H\\,\\I\\I\\ zone, $\\theta_{\\rm HII}$, is related to the mass loss rate from the WD as $\\dot M_{\\rm wind}^{-2}$. This implies that a relatively small decrease in the $\\dot M_{\\rm wind}$ can yield a large increase in the $\\theta_{\\rm HII}$ angle \\citep[see Fig.~1 of][]{cask12}. During the first days, when the $\\dot M_{\\rm wind}$ was at a maximum, $\\theta_{\\rm HII}/2 < i$ (i.e. the hot WD was below the H\\,\\I/H\\,\\I\\I\\ boundary), because we observed the warm pseudophotosphere. A small decrease in $\\dot M_{\\rm wind}$ could lead to opening of the H\\,\\I\\I\\ zone so that $\\theta_{\\rm HII}/2 > i$, i.e. the burning WD raised up the H\\,\\I/H\\,\\I\\I\\ horizon and thus we could observe directly its hot $\\lesssim 2.1$\\ro\\ photosphere at day 7.3. This could happen in a short time, because of a relatively low value of $i$. A transient increase of the \\textsl{EM} at day 7.3 supports this view.  \\subsection{Long-lasting super-Eddington luminosity}  Multiwavelength modelling the SED revealed that the luminosity of the WD was super-Eddington during the whole period of its nuclear burning (Table~2 here and Table~2 of the following paper).  Prior to the emergence of supersoft X-rays, only the lower limit of the WD's luminosity could be determined. During the presence of the two-temperature-type of the spectrum (i.e. for the first 4 days), it was given by the integral of directly measured fluxes (or their model) scaled with the distance. The model SED at day $\\sim 1$ and 3.8 corresponded to the luminosity of 6 and 2\\,$L_{\\rm Edd}$, respectively, for $M_{\\rm WD} = 1.3$\\mo. However, its true value was higher, because both the warm shell and the nebula converted only a fraction of the WD radiation into direction of the observer (Sect.~4.1.2, Eq.~(\\ref{eq:lbol}), Fig.~\\ref{fig:sketch}).  From day 7.3, measuring only the short Rayleigh-Jeans tail of the hot WD radiation, only the lower limit of its temperature and luminosity could be determined (see Appendix~A). Using the emission measure from the model SED and the $F_{\\rm h}(\\lambda)$ flux from the far-UV spectrum, allowed to determine the minimum temperature of the WD photosphere, $T_{\\rm h}^{\\rm min} \\sim 10^5$\\,K (Eq.~(\\ref{eq:tmin2}), Table~2) and the minimum of its luminosity, $L(T_{\\rm h}^{\\rm min}) = 13 - 4\\,L_{\\rm Edd}$ (Eq.~(\\ref{eq:lth})), at which the burning WD is just capable of producing the observed \\textsl{EM} during day 7.3 -- 19.5. Note that the luminosity of the nebular component itself corresponds to 1.24\\,$L_{\\rm Edd}$ at day 7.3 (Table~2).  The emergence of the first supersoft X-rays from the WD at day 26 allowed to determine its luminosity unambiguously to $L_{\\rm h} = 65\\pm 12$\\,$L_{\\rm Edd}$ (Sect.~3.5). Such a highly super-Eddington luminosity contradicts theoretical predictions, treating the nova event as a simple nuclear bomb explosion. According to \\cite{starrfield+08} the luminosity can exceed the Eddington limit only at the very beginning of the nova eruption with a rapid decline below it on a time-scale of hours.  However, the opposite cases were reported by some authors already during 70's and 80's decades of the last century. For example, \\cite{friedjung87} showed that the energy output of the nova FH~Ser was well above the Eddington limit for about two months after its optical maximum. The super-Eddington state was theoretically investigated by \\cite{shaviv01}, who considerred a reduction of the effective opacity due to the rise of a `porous layer' above the convective zone of the burning WD \\citep[see Fig.~1 of][]{shav+dot10}. These conditions enlarges the Eddington luminosity well above its classical value, calculated for the Thomson scattering opacity. The super-Eddington state also accelerates a thick continuum driven wind at high rates. The mass loss through the wind at a high rate of a few $\\times\\,10^{-4} - 10^{-5}$\\myr\\ during the hydrogen burning period of RS~Oph \\citep[][]{sk+08b} is consistent with that suggested for the super-Eddington state \\cite[Eq.~(30) of][for $v_{\\rm s}$ = 20 - 50\\kms]{shaviv01}. In addition, the presence of the bipolar jets for the first 30 days \\citep[][]{sk+08b,banerjee}, which suggests accretion throughout the disk at a high rate, also supports the super-Eddington state. For example, accretion at $\\dot m \\sim 2\\times 10^{-5}$\\myr\\ can release its binding energy $G M_{\\rm WD} \\dot m/R_{\\rm WD} \\sim 10^{39}$\\es, and thus help to sustain the high luminosity of the accretor for a long time.  Finally, it is important to note that the high energy output released during the outburst of RS~Oph, as given by the global model SED of this paper, is more than a factor of 10 larger than that so far derived. This requires a very high accretion rate, which is far above that can be obtained from a spherical wind mass loss of the giant \\citep[e.g.][]{schaefer09}. Considering that the primary source of the material powering the outburst is the wind from the giant, then a different, more efficient mechanism of the mass transfer onto the WD must be in the effect. This could be realized by the so-called wind Roche-lobe overflow (WRLOF), which can occur when the wind acceleration zone is stretched out to the Roche lobe of the giant \\citep[e.g.][]{moh+pod07,abate+13}. In this case the wind of the giant is focused towards the orbital plane and in particular towards the WD. The accretion rate predicted in the WRLOF regime can be 100 times larger than the rate expected from standard Bondi-Hoyle accretion \\citep[see][]{moh+pod07}.  More detailed discussion on the accretion mechanism is presented in the following paper, {\\em RS~Ophiuchi: The supersoft X-ray phase and beyond}, where other critical parameters are determined from the model SEDs during the SSS phase and following quiescence."
    },
    "1402/1402.3229_arXiv.txt": {
        "abstract": "{The embedded phase of star formation is the crucial phase where most of the stellar mass is assembled. Velocity-resolved spectra reveal an infalling envelope, bipolar outflows, and perhaps an infant circumstellar disk -- all locked together in a cosmic dance of gravitational collapse and magnetic winds. Densities and temperatures change by orders of magnitude as the protostar evolves, driving a chemistry as exotic as it is fascinating. I will review two examples of how to exploit chemistry and molecular spectroscopy to study the physics of low-mass star formation: energetic feedback and episodic accretion.}  \\FullConference{Frank N. Bash Symposium 2013: New Horizons in Astronomy (BASH 2013)\\\\ October 6-8, 2013\\\\ Austin, Texas }  \\begin{document} ",
        "introduction": "\\label{sec:intro} \\noindent For all of human history, we have wondered about the origins of the Earth, the Sun, and the rest of the solar system. The notion that stars form from the gravitational collapse of a molecular cloud core traces back to the ``nebular hypothesis'' developed by \\citet{swedenborg1734a}, \\citet{kant1755a}, and \\citet{laplace1796a}. Since then, many generations of astronomers have filled in the details of what such cloud cores look like, how they collapse, and how this results in a planetary system \\citep{cameron88a}. Others have studied the differences between low-mass and high-mass stars, between stars born in isolation or in clusters, and between star formation in the Milky Way or other galaxies \\citep{kennicutt12a}.  The state of affairs in the field of star formation has been reviewed periodically, in part through the \\emph{Protostars and Planets} series of conferences. The most recent edition, PPVI, took place in July 2013 in Heidelberg, Germany and featured chapters on molecular clouds \\citep{dobbs14a,andre14a}, embedded protostars \\citep{li14a,dunham14a}, chemical evolution \\citep{vandishoeck14a,ceccarelli14a}, and much more. Earlier reviews on various aspects of star formation were written by, e.g., \\citet{shu87a}, \\citet{vandishoeck93a}, \\citet{vandishoeck98a}, \\citet{lada99a}, \\citet{andre00a}, \\citet{klein07a}, \\citet{white07a}, and \\citet{mckee07a}.  The process of low-mass star formation is usually split into four to six steps, as illustrated in \\fig{sfpanels}. The first step (panel a) is a molecular cloud of typically a few pc in size with a mean density of about 100 \\pcc{} and a total mass of a few 1000 \\msun{} \\citep{bergin07b}. The cloud contains denser clumps of a few 100 \\msun{} each, which in turn contain individual dense cores (\\ten{4}--\\ten{5} \\pcc) of a few \\msun{} in mass and about 0.1 pc in size. In many star-forming regions, the clumps and cores line up in filaments \\citep{andre14a}.  \\begin{figure}[t!] \\centering \\includegraphics[width=\\textwidth]{starformation-en-green.eps} \\caption{Schematic view of isolated low-mass star formation, based on \\citet{shu87a}. Darker colors represent higher densities. The typical sizes and timescales for each step are indicated.} \\label{fig:sfpanels} \\end{figure}  Star formation properly starts when a dense core loses pressure support and begins to undergo gravitational collapse (\\fig{sfpanels}b). Due to the core's rotation, material falling in from the outer parts gets deflected towards the gravitational midplane. The young star soon launches a bipolar jet that carries away some of the core's angular momentum. The jet carves out a bipolar cavity from the infalling envelope and drives an outflow into the surrounding molecular cloud \\citep[\\fig{yso-parts};][]{bally07a,ray07a,arce07a,shang07a,frank14a}. Initially, when the young star is still deeply embedded in its natal core, it is called a Class 0 or Stage 0 protostar (\\citealt{andre93a}; see \\citealt{robitaille06a} for the distinction between observational classes and evolutionary stages). The transition to Stage I (\\fig{sfpanels}c) occurs when the star has accreted enough matter that its mass exceeds that of the remnant circumstellar reservoir \\citep{lada99a,robitaille06a}. The median lifetimes for Stage 0 and I protostars -- together called the embedded stage -- are 0.10 and 0.44 Myr \\citep{evans09a}.  After the envelope is cleared away, we are left with a pre--main-sequence T Tauri star surrounded by a circumstellar disk of gas and dust \\citep[\\fig{sfpanels}d;][]{bertout89a,dullemond07b}. Although disks of one form or another have now been detected in Stage 0 and I sources \\citep[e.g.,][]{jorgensen07a,takakuwa12a,tobin12a,murillo13a}, it remains unclear how and when disks first form and how big they grow in the embedded phase \\citep{li14a}. In T Tauri stars, disks can be anywhere from less than 10 AU to a few 100 AU in radius and their mass can be as high as 0.1 \\msun{} \\citep{andrews07b}. They last for a few Myr, during which the dust settles towards the midplane and grows first to larger grains, then to rocks and boulders, and ultimately to planets \\citep{johansen14a,raymond14a,helled14a,benz14a}. Meanwhile, the gas in the disk gradually gets dispersed into the interstellar medium \\citep{alexander14a} or accreted by giant planets. The combined effects of planet formation and gas dispersal can leave gaps and holes, producing a so-called transitional disk \\citep[\\fig{sfpanels}e;][]{espaillat14a}. After the gas is cleared away, the system turns into a debris disk in which rocky planets can continue to grow \\citep{matthews14a}. This phase can last for another 10-100 Myr before we end up with a main-sequence star and a mature planetary system (\\fig{sfpanels}f).  The current review covers two topics of recent interest in the embedded phase of star formation: energetic feedback (\\sect{feedback}) and episodic accretion (\\sect{epacc}). Both parts discuss in particular the use of chemistry and molecular spectroscopy as a probe of the the underlying physical processes. ",
        "conclusions": "\\label{sec:conc} \\noindent Despite decades of work, many aspects of low-mass star formation remain poorly understood. Recent observations with Spitzer and Herschel have shed new light on the energy balance in embedded protostars and on the relative importance of feedback from UV photons and shocks. Another topic of current interest is episodic accretion, i.e., a scenario in which a young star gains most of its mass in short bursts while spending most of the time in a state of low accretion. Chemistry and molecular spectroscopy are proving to be invaluable tools in the worldwide efforts to characterize and understand energetic feedback and episodic accretion in more detail."
    },
    "1402/1402.1476_arXiv.txt": {
        "abstract": "{ We re-examine the predictiveness of single-field inflationary models and discuss how an unknown UV completion can complicate determining inflationary model parameters from observations, even from precision measurements. Besides the usual naturalness issues associated with having a shallow inflationary potential, we describe another issue for inflation, namely, unknown UV physics modifies the running of Standard Model (SM) parameters and thereby introduces uncertainty into the potential inflationary predictions. We illustrate this point using the minimal Higgs Inflationary scenario, which is arguably the most predictive single-field model on the market, because its predictions for $A_\\ssS$, $r$ and $n_s$ are made using only one new free parameter beyond those measured in particle physics experiments, and run up to the inflationary regime. We find that this issue can already have observable effects. At the same time, this UV-parameter dependence in the Renormalization Group allows Higgs Inflation to occur (in principle) for a slightly larger range of Higgs masses. We  comment on the origin of the various UV scales that arise at large field values for the SM Higgs, clarifying cut off scale arguments by further developing the formalism of a non-linear realization of $\\rm SU_L(2) \\times U(1)$ in curved space. We discuss the interesting fact that, outside of Higgs Inflation, the effect of a non-minimal coupling to gravity, even in the SM, results in a non-linear EFT for the Higgs sector. Finally, we briefly comment on post BICEP2 attempts to modify the Higgs Inflation scenario.} ",
        "introduction": "Recently, the LHC has discovered a Higgs-like boson \\cite{Aad:2012tfa, Chatrchyan:2012ufa}, and Planck \\cite{Ade:2013uln} has reported precise measurements of the properties of the Cosmic Microwave Background (CMB).\\footnote{While this paper was in press, the even more exciting announcement of a measurement of $r$ was made by the BICEP2 collaboration \\cite{Ade:2014xna}, we briefly comment on this development in the context of Higgs Inflation in a note added in the conclusions.} In both cases, simplicity apparently rules. The LHC results are consistent with the Standard Model (SM), including the simplest linear realization of $\\rm SU_L(2) \\times U_Y(1)$ in the scalar sector, and rule out many exotic alternatives. The properties of the CMB as inferred by Planck, WMAP \\cite{Hinshaw:2012aka} and other ground based observations \\cite{Hou:2012xq,  Sievers:2013ica} are consistent with the Gaussian, adiabatic primordial curvature perturbations, typically predicted by single-field slow-roll models. This seemingly rules out many more exotic inflationary scenarios\\footnote{Although some of the apparently simplest scenarios, such as some power law single field models are also disfavoured by Planck data.}.  Both developments raise the stakes for the Higgs Inflation (HI) proposal \\cite{Bezrukov:2007ep}\\cite{DeSimone:2008ei,Bezrukov:2008ej} which aspires to use the SM Higgs boson as the single-field inflaton. The idea is to do so by adding the term $\\delta \\mathcal{L} = -\\xi (H^\\dagger H) \\, R$ to the combined Einstein-Hilbert and SM Lagrangians (where $H$ is the Higgs doublet and $R$ is the metric's Ricci curvature scalar), thereby making the Higgs sector into a non-minimally coupled inflationary model \\cite{Spokoiny:1984bd, Salopek:1988qh}. This seems a very benign, and arguably simple modification of known physics, since the new term is proportional to a dimensionless coupling ($\\xi$) that is allowed by the symmetries given the  SM field content.  At face value this model has many compelling features, no new fields are required beyond those describing particles now known to exist. Furthermore, it seems extremely predictive because all parameters except $\\xi$ are determined by non-cosmological physics, and $\\xi \\simeq 10^4$ is fixed by requiring the amplitude of primordial scalar fluctuations agree with CMB observations. Once this is arranged, the predictions for the scalar spectral index, $n_s$, and primordial tensor-to-scalar ratio, $r$, become parameter-independent at leading order (other than the dependence on SM parameters arising in the reheating analysis that is used to fix the number of inflationary $e$-foldings, $N_e$). Best yet, the predictions are successful: $n_s \\simeq 0.967$ and $r \\simeq 0.0031$ agree well with the Planck data.\\footnote{With the advent of BICEP2's measurement of $r = 0.20^{+ 0.07}_{-0.05}$ the later prediction is in conflict with the data. But, the Higgs inflation paradigm has since been modified {\\it post-hoc} to accommodate a larger $r$. See the comments at the end of the paper regarding this development.} See \\cite{Bezrukov:2013fka} for a recent review on this model.  This success, and the improved observational constraints, has led to a more systematic assessment of inflationary models in view of the observations \\cite{Martin:2013nzq}, with the HI model used as the benchmark model against which others are assessed in a Bayesian comparison. Indeed, such an analysis favours models for which inflation is not ruined by small parameter changes, and whose $n_s$ and $r$ predictions agree with the data as their parameters vary over a wide range of values. This tends to reward models with exponential potentials, like $V(\\phi) = A - B e^{-\\lambda \\phi}$, for which the slow-roll condition requires only that $\\phi$ be sufficiently large. This includes both the HI model and $R^2$ inflation \\cite{Starobinsky:1980te}\\footnote{See \\cite{Kehagias:2013mya} for a study of their essential equivalence in the large field regime.}. This result can also be viewed to be consistent with many models where exponentials arise in higher-dimensional theories, where the inflaton is a geometrical modulus (like the size, $r$, of an extra dimension) given that the associated energies can arise as powers of $1/r$ and the canonical field for such a quantity is $\\phi \\sim \\ln r$ \\cite{Burgess:2013sla, Roest:2013fha}. In particular, Ref \\cite{Burgess:2001vr} advocated extra dimensional models for exactly this exponential behavior far in advance of Planck data.  In this paper, we re-examine the predictiveness of single-field inflationary models, using the HI model as their poster child. We revisit the issue of the sensitivity of inflationary predictions to unknown UV physics, with the effects of this physics systematized within an Effective Field Theory (EFT) framework \\cite{Donoghue:1995cz, Burgess:2003jk}  for gravity. Beyond the `usual' UV sensitivity issues that are well known: the propensity of UV physics to ruin the flatness of the inflaton potential; and the sensitivity of slow-roll parameters to `Planck slop' --- {\\em i.e.} $1/M_p$ suppressed higher-dimension effective interactions, we identify another issue of UV sensitivity.\\footnote{Here $M_{p} = 2.44 \\times 10^{18} \\, {\\rm GeV}$ is the reduced Planck mass. We note that sensitivity to `Planck slop' is also called the $\\eta$-problem in some literature.}  The new issue we discuss first arises for inflationary models that are predictive in the sense that HI models are: that is, there are fewer free parameters than there are inflationary observables. In this case, the fact that the Renormalization Group (RG) running {\\em even at low energies} required to relate inflationary predictions to other observables is UV sensitive also introduces new parameters into the predictions for quantities like $n_s$ and $r$.\\footnote{The effect of UV physics, classified in terms of higher dimensional operators. modifying the running of the SM parameters was recently completely calculated for the first time in Ref.\\cite{Jenkins:2013zja}, for dimension six operators. We will use these results extensively in this paper in this application to cosmology.}  \\subsection{UV Issues}  We here briefly describe in more detail, and contrast, the various kinds of UV sensitivity that can arise, in order to set the context for the quantitative calculation in the next sections of their effects in the minimal HI model.  An EFT analysis of inflation leads to the well-known observation that UV physics generically tends to modify the inflaton potential so strongly that it ruins the flatness that is responsible for the slow roll. It typically does so because integrating out UV physics at a scale $M$ contributes to low-dimension operators in the EFT --- like corrections to the vacuum energy or scalar masses, $\\delta \\cL = - \\sqrt{-g} \\left( c_0 + c_2 \\, \\phi^2\\right)$ --- that are generically large: $c_0 \\propto M^4$ and $c_2 \\propto M^2$. This is the inflationary version of the standard `naturalness' problems that make it challenging to have light scalars within a generic EFT.  On the other hand, it is also known that once the low-dimension interactions are under control, UV physics can decouple from generic inflationary predictions \\cite{Kaloper:2002cs, Burgess:2003zw}, just like it does from other types of low-energy phenomena (provided the UV physics is adiabatic \\cite{Danielsson:2002kx, Martin:2000xs,Burgess:2002ub}). This is because corrections to high-dimension interactions are {\\em suppressed}, rather than enhanced, by the large scale. If $\\delta \\cL = - c_k \\, \\sqrt{-g} \\; \\phi^k$ then $c_k \\propto M^{4-k}$, which is suppressed for large $M$ if $k >4$. Of course the effective interactions satisfying $k \\le 4$ can still be problematic.\\footnote{In general the coefficient, $c_\\ssQ$, of an operator $Q$ in $\\delta \\cL$ varies as $M^{4-d_\\ssQ}$, where $d_\\ssQ$ is the full scaling dimension of $Q$ (including anomalous dimensions). Note that a sensitivity of large-field models to higher-dimensional Planck slop can be due to large field excursions generating large anomalous dimensions for operators that were initially suppressed, potentially spoiling inflation as it progresses. See however Ref.\\cite{Silverstein:2008sg} for a construction that avoids this.}  However even if such UV contributions are small in absolute size, they can still be large enough to ruin (or strongly perturb) inflation, since inflation requires not just that the inflaton mass be smaller than $M$; it must also be smaller than the Hubble scale, $\\cH \\sim V/M_p^2 \\ll M$. Because of this, interactions suppressed by powers of $1/M$ can still contribute non-negligibly to slow-roll parameters --- and so also to $r$ and $n_s$ --- even if they do not ruin inflation. For instance, a $c_6 \\, \\phi^6$ contribution to the potential competes with an $m^2 \\phi^2$ term whenever $c_6 \\, \\phi^4 \\propto \\phi^4/M^2 \\simeq m^2 \\lsim \\cH^2$. This can actually happen (even if $M \\simeq M_p$) because $\\cH$ is itself Planck-suppressed relative to the other scales in the potential. For most inflationary models, however, the slow-roll parameters are not predicted in terms of other observables, so the standard approach simply rolls all such UV contributions into the uncertainty in the values of the slow-roll parameters, allowing them to be ignored in practice.  Our focus in this paper is on a third way UV physics affects the low-energy inflationary model, distinct from the above two well-understood issues. It first arises when the inflationary model involves fewer parameters than there are inflationary observables, such as in Higgs Inflation. For the HI model, one measures the couplings within the scalar potential in particle physics experiments at comparatively low energy, and inflationary predictions are then made in terms of these parameters. This raises a technical complication because the field values, $(H^\\dagger H)_{\\rm inf} \\sim M_p^2/\\xi$, associated with inflation are enormous relative to those, $(H^\\dagger H)_{\\rm vac} \\sim v^2 \\simeq (246 \\; \\hbox{GeV})^2$, relevant to particle physics. The extrapolation of the potential to fields this large involves large logarithms, whose leading behaviour can be summed using Renormalization Group (RG) methods. This RG-based extrapolation is an important step when relating the large-field/high-energy inflationary potential to the small-field/low-energy parameters inferred from particle physics measurements at electroweak (EW) energies \\cite{DeSimone:2008ei, Bezrukov:2008ej, Bezrukov:2009db}.  Our main point is that this RG improvement of the potential is also sensitive to the details of a host of higher-dimension effective interactions, most of which are {\\em not} pure Higgs field operators. For instance, by contrast with $\\delta \\cL = - c_6 \\sqrt{-g} \\; \\phi^6$, an effective interaction like $\\delta \\cL = -\\frac14 \\, c_g \\sqrt{-g}\\; (H^\\dagger H) F_{\\mu\\nu} F^{\\mu\\nu}$ does not contribute at tree level to the scalar potential, because of the presence of the gauge fields. However it {\\em does} contribute at the quantum level because this operator contributes to the running of the corresponding gauge coupling of order $\\delta (1/g^2) \\sim c_g \\, m_h^2/16\\pi^2$ once the quantum fluctuations of the Higgs fields are calculated at one loop. This modification in the running of $1/g^2$ also feeds into the RG evolution of the other SM couplings at two loops, and this contributes to the running of the Higgs coupling, $\\lambda$ \\cite{Jenkins:2013zja}. As a consequence the value of the coupling $c_g$ can find its way into inflationary predictions.  Naively the size of any such contributions to $\\delta\\lambda$ would be expected to be very small. After all, if the effective interaction arises from integrating out a particle at mass $M$, then $c_g \\, m_h^2 \\propto m_h^2/M^2$, rapidly becomes very small for $M \\gg m_h \\simeq 125$ GeV.  Further, the specific example mentioned in the previous paragraph is  a two loop effect. However, there are also one loop effects of this form. Further, in the HI model, because inflation takes place at large values of the Higgs field, $H \\sim M_p/\\sqrt\\xi$, and $m_h$ is itself proportional to $H$, the effective Higgs mass can also be very large. Restricting to the contributions of operators of the form $\\delta \\cL \\sim H^2 F^2$, one finds a contribution to the running of $\\lambda$ {\\em at one loop} \\cite{Jenkins:2013zja} \\be \\label{d1grun} \\delta \\left( \\mu \\frac{d \\lambda}{d \\mu}\\right) \\supset \\frac{m_h^2}{16\\pi^2}\\left[ 9 \\, g_2^2 \\, C_{\\ssH \\ssW} + 3 \\, g_1^2 \\, C_{\\ssH \\ssB}  + 3 g_1 \\, g_2 \\, C_{\\ssH\\ssW\\ssB} \\right] \\sim \\frac{g^2 m_h^2}{16\\pi^2 \\, M^2} \\,, \\ee where the final relation indicates the order of magnitude with $g$ generically representing $g_1$ and $g_2$, the couplings of the $\\rm SU_\\ssL(2) \\times U_\\ssY(1)$ EW gauge bosons. See Ref  \\cite{Jenkins:2013zja} for details on the operator notation used here. This correction need not be inordinately small if $m_h \\sim M$ at the values of $H$ of interest, even if $M$ itself is very large.  There are principally two kinds of uncertainty in this kind of expression. The first is what value to use for the mass, $M$, of any new threshold. Part of the framework of minimal HI is the assumption that there are no new heavy particles beyond the SM between EW and inflationary energies, because any such thresholds generically could introduce effective couplings --- like $(H^\\dagger H)^3/M^2$, for example --- whose appearance within the potential could disturb the dynamics enough to destroy the inflationary slow roll.  However the mass, $M$, of the lowest new particle cannot be arbitrarily high. $M$ cannot be much larger than $ \\sim \\Lambda$, where $\\Lambda$ is the `unitarity scale', or the upper limit of the domain of validity of the semi-classical approximation \\cite{Burgess:2009ea}. $\\Lambda = \\Lambda(H^\\dagger H)$ is Higgs-field dependent \\cite{Bezrukov:2010jz}, and arises because the coupling to gravity is not renormalizable, and so the size of quantum effects can only be quantified within an EFT framework. Within this framework (as we review below) $\\Lambda \\sim M_p/\\xi$ for the small fields, $H \\ll M_p/\\xi$, relevant to particle physics; while $\\Lambda \\sim M_p/\\sqrt\\xi$ for the larger fields, $H \\simeq M_p/\\sqrt \\xi$, relevant to inflation. If we conservatively use $M \\sim \\Lambda \\propto M_p/\\sqrt\\xi \\sim H$ in the inflationary regime, and that $m_h \\propto H$ there, we see that the correction in Eq.~\\pref{d1grun} can be comparable to one-loop contributions computed within the SM. Effects such as this can be large enough to visibly change the implications for HI in the $n_s - r$ plane, as is illustrated in Fig.~\\ref{Fig;nsvsr}.   \\begin{figure}[t] \\centering \\includegraphics[width=5in,height=3.5in]{plot1.pdf} \\caption{The potential effect of the unknown UV completion on the expectation for CMB parameters in the Higgs inflation scenario. The red dot is the prediction in HI for the scalar to tensor ratio $r$, and the spectral index $n_s$, without the effect of higher dimensional operators modifying the RG running. See Section \\ref{nonlinear} for the details of how this prediction is obtained. The black line is the span of expected values for these parameters when the higher dimensional operators are also marginalized over. (The thickness of the line in the direction of $r$ is exaggerated so that the line is visible.) The figure also shows the one and two sigma regions of Fig 4 of Ref. \\cite{Ade:2013uln}. The larger red regions are Planck and WMAP data + BAO + $\\rm \\Lambda CDM$ + r allowing running of $d n_s/d k$. The smaller blue regions are Planck and WMAP data + BAO + $\\rm \\Lambda CDM$ + r not allowing a running of $d n_s/d k$.} \\label{Fig;nsvsr} \\end{figure}  The other uncertainty in these estimates is whether or not all other higher dimensional operators in the non-renormalizable EFT actually vanish. They do not for any known proposals for weakly coupled physics beyond $\\Lambda$ (such as string theory, or higher-dimensional gravity, for example). All of the higher dimensional operators will be generated by renormalization, so any vanishing of all these terms, if accomplished, will necessarily only occur at one scale. Further, one need not consider this question to be an exotic one purely in the context of gravity. We discuss in Section \\ref{unitaritysection} how attempts to banish these operators can be mapped to analogous statements on unitarity violation involving massive spin one states in an EFT, with no need to invoke gravity. Attempts to argue away these operators would in this manner have broader implications for our understanding of unitarity violation and renormalization in many EFTs.  But these strong arguments, and the  absence of examples, does not remove the logical possibility that such physics might exist; sufficiently suppressing {\\em all} dangerous dimension-six interactions at the scale $\\Lambda$. It is difficult to say more without a specific and precise proposal for what the UV physics is that must enter at scale $\\Lambda$,what the coefficients of the operators will be\\footnote{See for instance Refs. \\cite{Antoniadis:2009rn, Antoniadis:2010nb} for a classification of the operators that might appear in the context of the MSSM. See also Ref. \\cite{Chatterjee:2011qr} for a study of Higgs inflation embedded in the MSSM. We also note that some simple models have also been proposed to UV complete Higgs inflation and avoid the unitarity bound,see for example \\cite{Giudice:2010ka}.}. Given our current lack of knowledge of physics beyond the Standard Model, our own point of view is that these unknown order-unity coefficients are likely to be nonzero and so represent intrinsic theoretical uncertainties that must be propagated through to low energy observable quantities in EFT's, such as CMB observables in HI\\footnote{In a similar spirit, see \\cite{Branchina:2013jra} for a discussion of the sensitivity of the stability of the EW vacuum to new physics.}.  The importance of these threshold-like terms within RG equations was recently emphasized for non-cosmological applications in Ref.\\cite{Jenkins:2013zja}, as part of a systematic renormalization program of the SM EFT (with full flavour structure) completely carried out in Refs. \\cite{Grojean:2013kd,Jenkins:2013zja,Jenkins:2013wua,Alonso:2013hga}. In Ref.\\cite{Jenkins:2013zja} the complete modification of the running of the parameters present in the renormalizable SM Lagrangian due to dimension six operators was explicitly calculated.  In what follows we use these results to illustrate how the running of SM parameters can be modified in the case of HI. We discuss how this impacts attempts to predict $n_s$ and $r$ in this model, and derive the results illustrated in Fig.~\\ref{Fig;nsvsr}.  Our broader lesson is this: although we discuss in detail HI, similar issues should arise within the SM RG in any attempts to link EW scale physics with the higher scales involved in inflationary (and other cosmological) scenarios\\footnote{Any complete account of reheating into some sector that contains the Standard Model seems to force this issue upon us by directly or indirectly coupling the SM degrees of freedom to the inflaton, for example.}. UV sensitivity is a many-headed hydra, and it is only with the development of more predictive models that this latest version has become potentially relevant.  The outline of this paper is as follows, in Section II we discuss HI and the cut off scales present in theories of this form. In Section III, we discuss the RG evolution used in these theories, and we outline the contributions to the RG equations that we include that were previously neglected. We then demonstrate how these corrections impact predictions in these theories based on EW scale measurements. Finally, in Section IV, we conclude. ",
        "conclusions": "We have re-examined the issues of UV sensitivity in inflationary single field models, focusing on the interesting case of Higgs Inflation. The effect of unknown higher dimensional operators were shown to have an observable impact on CMB predictions in this case. This is an irreducible theoretical uncertainty (in our view) until the exact UV that completes the theory is specified. It is not sufficient to banish higher dimensional operators that are composed only of Higgs fields in models of this form to maintain predictivity. The higher dimensional interactions of the same dimension extensively mix, at sufficient loop order. Further the higher dimensional operators mix down and modify the SM parameter running in a manner that depends on the background field value. This introduces UV sensitivity at the scale $M_{p}/\\xi$ through the RG equations, in an interesting manner. The requirement of an exponentially flat potential makes some inflationary models particularly sensitive to these effects.  Note added on recent developments: The recent BICEP2 measurement \\cite{Ade:2014xna} of a comparatively large primordial tensor fluctuation, $r = 0.20^{+ 0.07}_{-0.05}$, puts some pressure on the Higgs Inflationary scenario which predicts smaller $r$ for inflation driven by the exponential rollout from the asymptotically constant Einstein-frame potential at large fields. (See Ref. \\cite{Cook:2014dga}, for example, for a recent discussion). Other recent works \\cite{Hamada:2014iga, Bezrukov:2014bra} counter this with ways to evade the problem in special parts of parameter space. (For instance one can choose special values for $m_t$ and $m_h$ --- though not within the one-sigma measured values --- such that the critical point in the SM Higgs potential occurs at scales similar to those required by the BICEP2 measurements.) Once this is done a larger value of $r$ can be obtained, potentially consistent with the BICEP2 results.  While we are willing to take these claims at face value, we would make the following comment: the larger value of $r$ so obtained comes at the expense of a much smaller value of $\\xi$: $\\xi \\sim 10$ rather than $\\sim 10^4$. This is worrisome for the control of approximations used, the point initially raised in Ref.~\\cite{Burgess:2009ea}, since it is precisely the large value of $\\xi$ that provides the hierarchy between the Planck scale $M_p$, the large-field unitarity scale, $\\Lambda \\sim M_p/\\sqrt\\xi$, and the inflationary Hubble scale $H \\sim M_p/\\xi$. For $\\xi$ of order 10 the unitarity scale is only 3 times larger than the Hubble scale during inflation, and both are uncomfortably close to the Planck scale. The effects of higher-dimension operators emphasized in this article are also a concern in this case; with the range of predicted values for both $r$ and $n_s$ being much larger than their measured errors. In general, smaller $\\xi$ implies close proximity to the UV `Planck wall', thereby sharpening all issues associated with the unknown UV completion at these scales."
    },
    "1402/1402.7114.txt": {
        "abstract": "{ We review recent progress on cosmological issues and theoretical properties of modified gravity theories. In particular, we explicitly explore the conformal transformation, the Starobinsky inflation, and a unified scenario of inflation and late time acceleration in $F(R)$ gravity and $F(T)$ gravity (extended teleparallel gravity). Furthermore, we examine neutron stars and the hyperon problem in $F(R)$ gravity. Moreover, for loop quantum cosmology (LQC), the natures of finite-time future singularities in $F(T)$ gravity are presented. In addition, we investigate $F(T)$ gravity theories from the Kaluza-Klein (KK) and Randall-Sundrum (RS) theories. } %%%%%%%%%%%%%%%%%%%%%  \\title{Universe acceleration in modified gravities: $F(R)$ and $F(T)$ cases} \\FullConference{KMI International Symposium 2013 on \u0081gQuest for the Origin of Particles and the Universe\",\\\\ %``Quest for the Origin of Particles and the Universe'' 11-13 December, 2013\\\\ Nagoya University, Japan} \\ShortTitle{Universe acceleration in modified gravities}  \\begin{document}   % % This modifies the name of the authors % in the headers of the pages % \\makeatletter \\setbox\\@firstaubox\\hbox{\\small K. Bamba, S. D. Odintsov} \\makeatother   %%%%%%%%%%%%%%%%%%%%%%%%%%% %%%  Section 1 %%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "Quite precise cosmological observations such as Type Ia Supernovae have revealed the current accelerated expansion of the universe. For the homogeneous and isotropic universe, there are two representative ways of explaining this phenomenon. First is the introduction of dark energy with negative pressure in general relativity. Second is the modification of gravity on the long distance, e.g., $F(R)$ gravity~\\cite{Capozziello:2002rd, Carroll:2003wy, Nojiri:2003ft} (for reviews on dark energy and modified gravity, see, for instance,~\\cite{Nojiri:2010wj, Nojiri:2006ri, Capozziello:2011et, Bamba:2012cp}).  In this article, among various modified gravity theories, we concentrate on $F(R)$ gravity and $F(T)$ gravity with $T$ the torsion scalar in teleparallelism~\\cite{A-P} and review the recent developments on cosmological issues and their theoretical properties. %%%%% %%% Unit %%% We use units of $k_\\mathrm{B} = c = \\hbar = 1$ and denote the gravitational constant $8 \\pi G$ by ${\\kappa}^2 \\equiv 8\\pi/{M_{\\mathrm{Pl}}}^2$ with the Planck mass of $M_{\\mathrm{Pl}} = G^{-1/2} = 1.2 \\times 10^{19}$\\,\\,GeV. %%%%%%%%%%%%  %%%%% Structure %%%%% The article is organized as follows. %%%%%%%% In Sec.\\ II, we review various cosmological subjects in $F(R)$ gravity. First, we explain the action and derive the gravitational equations. We explore the conformal transformation from the Jordan frame to the Einstein frame and its inverse transformation. After that, we investigate $R^2$ inflation, namely, the so-called Starobinsky inflation. Also, we present a unified model of inflation and dark energy era, and describe the properties of the finite-time future singularities. Moreover, we study neutron stars and the issue of hyperon. For the recent study of stars in dilaton gravity, see~\\cite{Fiziev:2014eba}. %%%%%%%% Next, in Sec.\\ III, we explain significant cosmological issues in $F(T)$ gravity. To begin with, we present the formulation of teleparallelism. We state the impossibility of the conformal transformation from $F(T)$ gravity in the Jordan frame into pure teleparallel gravity with a scalar field, which corresponds to the Einstein frame in the ordinary curvature gravity. In addition, we examine the analogue of the Starobinsky inflation, i.e., $T^2$ inflation. Furthermore, we illustrate the unification of inflation in the early universe and the late time cosmic acceleration at the dark energy dominated stage. We also mention the finite-time future singularities in loop quantum cosmology (LQC). Moreover, we study a conformally invariant $F(T)$ gravity theory. We further explore trace-anomaly driven inflation. In Sec.\\ IV, we review our main results on $F(T)$ gravity theories from the Kaluza-Klein (KK) and Randall-Sundrum (RS) theories in Ref.~\\cite{Bamba:2013fta}. %%%%%%%% In Sec.\\ V, summary are described. %%%%%%%%%%%%%%%%%%%%%      %%%%%%%%%%%%%%%%%%%%%%%%%%% %%%  Section 2 %%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "In the present article, we have reviewed recent progress on issues of cosmic acceleration and theoretical natures of $F(R)$ gravity and $F(T)$ gravity. %%%%% In the former part, we have described various cosmological and theoretical problems in $F(R)$ gravity. We have examined the conformal transformation from $F(R)$ gravity (i.e., the Jordan frame) to general relativity with a scalar field (i.e., the Einstein frame) and its inverse version. In addition, we have studied $R^2$ inflation (the Starobinsky inflation). Also, we have explained a unified scenario of inflation and the dark energy dominated stage. Moreover, we have presented the classification of the finite-time future singularities. Furthermore, we have examined neutron stars and the hyperon issue. %%%%% Next, in the latter part, we have stated a number of cosmological issues as well as theoretical properties in $F(T)$ gravity. First of all, we have written the basic formulation of teleparallelism. Then, we have remarked the impossibility of the conformal transformation from the Jordan frame to the Einstein frame. We have further investigated the analogue of the Starobinsky inflation, namely, $T^2$ inflation. In addition, we have demonstrated the unification between inflation and the late time accelerated expansion of the universe. Furthermore, we have mentioned the finite-time future singularities in loop quantum cosmology (LQC). Moreover, we have shown a conformally invariant $F(T)$ gravity theory. Finally, we have reviewed our main results of Ref.~\\cite{Bamba:2013fta} in terms of derivations of $F(T)$ gravity theories from the Kaluza-Klein (KK) and Randall-Sundrum (RS) theories. %%%%%%%%  It could be expected that through the investigations on various theoretical aspects of a number of gravity theories, we can acquire some clues to resolve the issues of cosmic acceleration including the mechanism of inflation in the early universe and the origin of dark energy at the present time.      %%%%%%%%%%%%%%%%%%%%%%%% %%% Acknowledgments %%%%%%%%%%%%%%%%%%%%%%%% %"
    },
    "1402/1402.2369_arXiv.txt": {
        "abstract": "{We report a multi-wavelength study of four new planetary nebula (PN) candidates selected from the INT/WFC Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS) and Deep Sky Hunter (DSH) catalogues. We present mid-resolution optical spectra of these PNs. The PN status of our sample was confirmed by optical narrow-band images and mid-resolution spectra. Based on the locations of these objects in the log (H$\\alpha$/[\\ion{N}{II}]) versus log (H$\\alpha$/[\\ion{S}{II}]) diagnostic diagram, we conclude that these sources are evolved low-excitation PNs. The optical and infrared appearances of these newly discovered PNs are discussed. Three of the new nebulae studied here are detected in infrared and have low infrared-to-radio flux ratios, probably suggesting that they are evolved. Furthermore, we derive the dynamical ages and distances of these nebulae and study the spectral energy distribution for one of them with extensive infrared archival data. } ",
        "introduction": "Planetary nebulae (PNs) are the evolutionary end products of low- and intermediate-mass (1 $\\sim$ 8 M$_\\odot$) stars. The number of PNs detected in the Milky Way is about 2700, although the total population of PN is believed to be about one order of magnitude higher. The list of detected PNs includes $\\sim$1500 in the Galactic PN Catalogue of  \\citet{Kohoutek01} and $\\sim$1200 PNs listed in the Macquarie-AAO-Strasbourg H$\\alpha$ Catalogue of Galactic PNs (MASH; Parker et al. 2006; Miszalski et al. 2008). The discrepancy between the detected and expected numbers of PNs in the Galaxy is believed to be primarily due to dust extinction, as most of the PNs are in the Galactic plane. Many evolved PNs have low surface brightness and can be easily missed in the optical surveys.  In this paper, we present a study of four highly obscured PN candidates that are located within both regions of the INT/WFC Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS; Drew et al. 2005)/Deep Sky Hunter (DSH) and the {\\it Spitzer} Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE)/{\\it Wide-field Infrared Survey Explorer} (WISE) infrared surveys. The sample is selected from a catalogue of 781 new PN candidates observed by IPHAS \\citep{Viironen09a} and a list of 36 un-catalogued Galactic nebulae observed from {\\it Digitized Sky Survey} (DSS) and {\\it Two Micron All Sky Survey} (2MASS) surveys \\citep{Kronberger06}. Almost all of these objects are located close to the Galactic plane and are very low-brightness nebulae, indicating that they may have high dynamic ages or/and suffer larger interstellar extinction.  One of the scientific goals of the IPHAS is searching for PNs with low surface brightness or faint extended nebulae \\citep{Mampaso06, Sabin10}. The IPHAS has mapped all northern Galactic plane coverage to $\\vert$b$\\vert\\leqslant$5$^\\circ$ by three filters using the Wide Field Camera (WFC) on the 2.5 m Isaac Newton Telescope (INT) at the Observatorio del Roque de los Muchachos (La Palma, Spain) \\citep{Drew05}. The spatial resolution of each image is 0.$\\arcsec$33 and the total area of the survey covers  $\\sim$1800 deg$^{2}$. Two broadband Sloan filters (r$\\arcmin$ and i$\\arcmin$) and one narrow-band H$\\alpha$ filter ($\\lambda_{c}$ = 6568 \\AA, $\\Delta\\lambda$ = 95 \\AA) were used with corresponding exposure times of 30, 10, and 120 s, respectively. The measured magnitude range for point sources in the survey is $r\\arcmin\\sim$ 13 to 22 mag. The survey provides us an opportunity to obtain better estimates of Galactic PN population, the PN distribution density, the PN formation rate, and the chemical abundance of the Galaxy. Until now, a number of new PNs \\citep{Viironen09b, Sabin10, Sabin11} have been discovered based on results from this survey.  The DSH project is a group involving many amateur astronomers with the goal of discovering new nebulae by searching optical and near-infrared archival data \\citep{Kronberger06, Jacoby10}.  Although the ionized structures of PNs are well determined by the optical observations, many faint evolved Galactic PNs suffer from high extinction, making them difficult to detect in the optical wavelengths. Unlike optical observations, infrared observations are hardly affected by dust extinction and therefore can much more effectively detect PNs. The recent infrared observations of {\\it Spitzer} GLIMPSE \\citep{Benjamin03, Churchwell09} and the WISE \\citep{Wright10} have made possible new ways to study Galactic PNs (e.g. Cohen et al. 2005). These results show that the optical appearances of PNs and their infrared counterparts may differ, and the infrared images can also help us to understand the distributions of dust components.  The main aim of this present work is to identify the nature of four new IPHAS/DSH PNs by the narrow-band imaging and follow-up spectroscopic observations. The observations and data reductions are described in Section \\ref{s2}. In Section \\ref{s3}, we present the results of imaging and spectroscopy in the visible and mid-infrared for these objects. The log (H$\\alpha$/[\\ion{N}{II}]) versus log (H$\\alpha$/[\\ion{S}{II}]) diagnostic diagram is discussed in Section \\ref{s4}. In Section \\ref{s5}, we present the study of spectral energy distribution for one of the new PNs. The dynamical ages and distances of these nebulae are discussed in Section \\ref{s6}. Finally, the conclusions are given in Section \\ref{s7}. ",
        "conclusions": "\\label{s7}  Recent H$\\alpha$ surveys with high spatial resolutions and sensitivities have provided us a powerful tools to study PNs obscured in the Galactic plane. More than one thousand new PNs and PN candidates have been discovered in the MASH images and IPHAS photometric catalogues \\citep{Mampaso06, Parker06, Miszalski08, Viironen09b, Corradi12}, containing examples of young, compact PNs and old, extended, faint PNs. The nature of some of them is still unknown and requires spectroscopy. Here, we present an optical and infrared study of a sample of four new PN candidates discovered in the IPHAS and DSH catalogues. We have confirmed the PN status of these objects by the optical narrow-band images and mid-resolution spectra. From the locations of the new PNs in the S2N2 emission-line-ratio diagnostic diagram, we find that the nebulae are low excitation PNs with high dynamic ages.  Three new PNs in our sample (PN1, PN2, and PN3) are found to have obvious infrared counterparts. Some of them (PN2 and PN3) display prominent infrared emissions in the central regions, suggesting that most dust masses of these objects are located in the central parts. The infrared-to-radio flux ratios of these three PNs (PN1, PN2, and PN3) are generally lower than those of other PNs in the GLIMPSE II and GLIMPSE 3D fields, suggesting that they are more evolved. We constructed the SED of PN3 by collecting the photometries from a variety of astronomical data sets, which cover the wavelength range from visible to far-infrared. The observed infrared fluxes are successfully fitted by two dust components with temperatures of 231 K and 60 K, respectively. This PN might have a dusty envelope, which heavily obscures its optical counterpart. We estimated the distances and ages of these new PNs. They are relatively old with an averaged age of $\\sim$ 8900$\\pm$4500 years, which is consistent with the results obtained from the S2N2 diagnostic diagram and the infrared-to-radio ratios."
    },
    "1402/1402.0700_arXiv.txt": {
        "abstract": "The problem of modeling the distribution of dark matter in galaxies in terms of equilibrium configurations of collisionless self-gravitating quantum particles is considered. We first summarize the pioneering model of a Newtonian self-gravitating Fermi gas in thermodynamic equilibrium developed by Ruffini and Stella (1983), which is shown to be the generalization of the King model for fermions. We further review the extension of the former model developed by Gao, Merafina and Ruffini (1990), done for any degree of fermion degeneracy at the center ($\\theta_0$), within general relativity. Finally, we present here for the first time the solutions of the density profiles and rotation curves corresponding to the Gao et. al. model, which have a definite mass $M_h$ and circular velocity $v_h$, at the halo radius $r_h$ of the configurations, typical of spiral galaxies. This treatment allow us to determine a novel core-halo morphology for the dark matter profiles, as well as a novel particle mass bound associated with them. ",
        "introduction": "One of the most quoted papers in the study of dark matter distribution in galactic halos is certainly the work of Tremaine \\& Gunn \\cite{PRL79Tremaine}, there the authors established a lower limit on the mass of neutrinos composing galactic halos, by considering an isothermal classical distribution of self-gravitating particles, and imposing quantum constraints on the phase space density in the core of the galaxies. This treatment presenting a peculiar mixture of quantum and classical considerations has always attracted the attention and the suspicion of many astrophysicists and theoretical physicists. It was in particular the opinion of one of us (R. Ruffini), that a self-consistent treatment of quantum constraints in a self-consistent quantum description of the microphysical system was needed. A long effort so started. ",
        "conclusions": "We have recently returned to the Gao et. al. work, and propose a completely different way for solving the boundary condition problem for the system (\\ref{eq:1}--\\ref{eq:2}), in order to fulfill the observationally inferred values of typical dark matter halos in spiral galaxies as given in \\cite{AJ08deBlok}. Namely, for given initial conditions $M_0=\\nu_0=0$, arbitrary $\\theta_0$ (depending on the chosen central degeneracy), and defining the halo radius $r_h$ at the onset of the flat rotation curve, we solve an eigenvalue problem for the central temperature parameter $\\beta_0$, until the observed halo circular velocity $v_h$ is obtained. After this, we solve a second eigenvalue problem for the particle mass $m$ until the \\textit{observed} halo mass $M_h$ is reached at the radius $r_h$.  The quest has been to use all these information in order to put constraints on the mass of the `ino' in galactic halos by introducing the observational properties possibly to be utilized in this research.  Interestingly enough, as detailed in \\cite{PRL13Ruffini}, it turns out that only for an specific range of $\\theta_0>0$ these two eigenvalue problems can be solved together, implying as a consequence, a novel reach morphology for the density profiles as well as a novel particle mass bound associated with it. The density profiles presents a quantum degenerate core, followed by a low degenerate plateau until it reach the $r^{-2}$ Boltzmannian regime corresponding to the flat part in the rotation curve. In Fig.~\\ref{fig:4} we show a family of density profiles for different values of $\\theta_0$ which fulfills the mentioned halo constraints. We also plotted for comparison the purely Boltzmannian curve which agrees with the same \\textit{observed} halo magnitudes.  \\begin{figure}[!hbtp] \\centering \\includegraphics[width=1.\\hsize,clip]{rhoQsNKorea.eps} \\caption{Physical density profiles for specific ino masses $m$ and central degeneracies $\\theta_0$ fulfilling the observational constraints $M_h=1.6\\times10^{11} M_\\odot$ and $v_h=168$ km/s at $r_h=25$ Kpc (as taken from \\cite{AJ08deBlok} and detailed in \\cite{PRL13Ruffini}). In dot-dashed line the purely Boltzmannian profile for comparison.}\\label{fig:4} \\end{figure}  As can be seen from Fig.~\\ref{fig:4}, we obtain from this novel analysis a more stringent lower mass bound for the `ino' mass $m$, which is $\\sim10$ times higher than the ones inferred in \\cite{AA83Ruffini} and \\cite{PRL79Tremaine}. This is, $m\\geq0.42$ keV$/c^2$ for typical spiral galaxies.  It is interesting that the quantum and relativistic treatment of the configurations considered here are characterized by the presence of central cored structures unlike the typical cuspy configurations obtained from a classic non-relativistic approximation, such as the ones of numerical N-body simulations in \\cite{ApJ97Navarro}. This naturally leads to a possible solution to the well-known core-cusp discrepancy \\cite{AJ01deBlok}."
    },
    "1402/1402.2533_arXiv.txt": {
        "abstract": "{ Gamma-rays from cosmological sources contain information about gamma-ray interactions. Standard model and non-standard model photon interactions along the path between the source and the observer can lead to changes in the energy or state of the photons, which in turn alters the observed energy spectrum of the source. In general, these interactions are a function of photon energy as well as source distance.  Here we show how existing high energy gamma-ray observations of blazars can be used to constrain the coupling of axion-like-particles (ALPs) to the photon.  The same ALP-photon coupling that has been invoked to explain the observations of TeV blazars beyond their pair production horizon is shown to have an effect of the data set of  \\textit{Fermi} blazars.} ",
        "introduction": "Cosmological gamma-rays are an ideal probe of relatively rare photon interactions. In addition to expected standard model interactions, like those with the extragalactic background light (EBL) through pair production, any additional interaction which alters the energy or state of the photon can in principle be probed.  The process of constraining gamma-ray interactions beyond the standard model depends on knowledge of the source gamma-ray spectra, the quality of rapidly growing data sets of high energy gamma-ray observations, and a sufficiently well understood modeling of the EBL. As a test case, we consider a group of gamma-ray sources at approximately the same distance, using the  high redhift ($z \\sim 1$) portion of the {\\textit Fermi} blazar catalogue \\cite{oai:arXiv.org:1108.1435} as an example. In addition, we discuss a similar set of TeV blazars clustered at $z\\sim 0.1$ (see references in \\cite{Reesman:2013jra}) . We expect blazars at the same distance to have their spectra be attenuated in a similar way by the EBL (for a recent review see \\cite{Costamante:2013sva}).  If we assume a well-determined model for the EBL then, in principle,  a comparison of the observed spectra to the spectra at the source can probe new photon couplings. This requires some assumptions about the intrinsic spectra, namely that they are a power  law. In practice, we then use the residuals between the observed (as measured by the {\\it Fermi} blazar data set) and expected EBL attenuation as an example of the range of constraints that can be placed on rare photon interactions. There have been observations of distant gamma-ray sources which suggest that the Universe is less opaque to gamma-rays than expected \\cite{Aharonian:2005gh, Aleksic:2011hr, Acciari:2009zy, Aliu:2008rg}. Rare photon couplings have recently been proposed to explain blazar spectra \\cite{SanchezConde:2009wu, Wouters:2013iya, Tavecchio:2012um, Wouters:2012qd}, leading  to general questions about novel couplings of cosmological photons, which we show could be probed by the \\textit{Fermi} blazars.  Specifically, we examine  the mixing of the photon with a pseudoscalar axion-like particle (ALP).  ALPs couple to photons via a two-photon vertex, just as the axion\\cite{PhysRevLett.38.1440}  does.  Unlike the axion, which has a definite relationship between this coupling and its mass, the ALP's mass is independent of coupling \\cite{Kim:1998va, Mirizzi:2006zy}.  In this paper we explore the effect of ALPS on the propagation of cosmological gamma-rays (see also \\cite{DeAngelis:2011id,DeAngelis:2007yu,SanchezConde:2009wu}). In Sec.~\\ref{sec:blazar} we will show how ALP-photon oscillations manifest in the observed energy spectra, in Sec.~\\ref{sec:calc} we discuss the data sets and proposed analysis technique, and finally we conclude with our results in Sec.~\\ref{sec:results}. Averaging over the stochastic intergalactic magnetic field domain structure, we find that the \\textit{Fermi} sources can exclude couplings larger than  10$^{-11}$ GeV$^{-1}$ for ALP masses below 10$^{-4} \\mu$eV. ",
        "conclusions": "\\label{sec:results} We have shown how to probe the scale of ALPs coupling to high energy photons if we assume that the residual between the observed optical depth and a best fit model (\\textit{Fermi} collaboration \\cite{Ackermann:2012sza}), cannot be exceeded by the ALP-photon interconversion associated with this coupling. For a detailed calculation of strong mixing in the IGMF see the DARMA model in \\cite{DeAngelis:2011id,DeAngelis:2007yu}.  In Fig.~\\ref{fig:constraint} we illustrate the results of this method.  As in Fig.~\\ref{fig:paramspace}, the slanted region corresponds to the paramter space such that $E_{crit}=30$ GeV and  the spread is for a range of IGMF $B$ field values. The actual calculation was done for $B=1$ nG, causing the maximal allowed ALPS coupling to lie on the right side of the 30 GeV band. Had we used a smaller magnetic field value (for example, $B=0.1$ nG) then the upper-bound of the mass would be a factor of three smaller for the same E$_{crit}$. The two constraint points shown correspond to the exact limit from the \\textit{Fermi} data assuming a critical energy of 30 GeV.  Moving to lower $m_{\\mu eV} $ corresponds to lower values of E$_{crit}$, which would still allow all the photons from the source to oscillate. Therefore points to the left of the shown constraints result in the same ALP distortion.  Similarly, points below the data constraints and along the lines of constant E$_{crit}$ correspond to longer decay lengths (lower probabilities for conversion) and are therefore allowed. The green shaded region is excluded by this data, under our assumptions. The red shaded region corresponds to points that are acceptable for \\textit{enhancement}, but unacceptable for \\textit{extra attenuation}. Both solutions correspond to roughly the same coupling, with the \\textit{enhancement} solution for slightly larger coupling. \\footnote{Back conversion in the Milky Way was used to constrain ALP-photon coupling using the lack of 100 MeV gamma rays from SN87a \\cite{Simet:2007sa}.  In this analysis, high energy ALPS are produced in the hot core of the supernova and then back convert in the Galaxy's magnetic field.  The absence of high energy gamma rays from SN87a then constrains $g_{11} \\geq 1$ for $m_{\\mu eV} <  10^{-3}$.  For the parameter range considered here, back conversion of blazar ALPs to photons would also occur in our Galaxy, but at a level minimal compared to the amplitude of residuals we consider.}    \\begin{figure} \\centering \\includegraphics[width=0.65\\textwidth]{ALP_constraint_update.eps} \\caption{ALP photon coupling and ALP mass  as constrained by HE blazar observations made by the Fermi collaboration. CAST line from \\cite{Andriamonje:2007ew}.} \\label{fig:constraint} \\end{figure}  Historically, the effect of ALPs on blazars has focused on examining the spectrum of an individual source and attributing a break in the assumed source power-law to ALPs. These analyses run Monte Carlo simulations over ALP parameter space, including a parameterization of magnetic field domains, among other quantities \\cite{ Tavecchio:2012um, SanchezConde:2009wu}. Our method instead illustrates the range of constraints possible if we integrate over the spectral variations of many sources, at approximately the same distance, at the same time. This is completely analogous to the EBL measurement made by \\textit{Fermi} and hence why we can use their result.  In a perfect world, improving the data set for TeV blazars at $z\\sim 0.1$ could be used in concert with the \\textit{Fermi} blazars to better probe slightly larger ALP masses at roughly the same coupling, in particular in the region where both experiments are measuring the same EBL \\cite{Reesman:2013jra}.  In \\cite{SanchezConde:2009wu}, ALP mixing in both the source and the IGMF was considered for 3C 279 (z$=$0.536) and PKS 2155-304 (z$=$0.117). Their fiducial model used the values g$_{11}=8.77$ and m$_{\\mu eV}=10^{-4}$. ALP-photon mixing can also be used to reduce the photon opacity within the source itself -- in a recent paper \\cite{Tavecchio:2012um} the possibility of ALP-photon conversion in the source PKS 1222+216  (z$=$0.432) was explored as a possible explanation of reducing the source opacity. They considered the presence of a strong B field in/around the source that allowed for $\\gamma\\rightarrow a$ conversion such that the gamma-ray could escape the broad line region of the blazar. Once escaped, the ALP could convert back $a\\rightarrow \\gamma$ in either the host galaxy, or in the IGMF. In order to explain both the HE and VHE observations (from \\textit{Fermi} and MAGIC) using a standard blazar model, the authors needed to invoke ALPs with an inverse coupling g$_{11}=$1.4 in a (source) magnetic field $B_{G}=0.2$.   That is to say, if ALPs are required to extend the TeV horizon  and/or reduce the source opacity of several 100 GeV photons in blazars, we would expect correlated ALP effects, at a scale demonstrated in this paper, in the set of \\textit{Fermi} blazars clustered near redshift 1."
    },
    "1402/1402.5914_arXiv.txt": {
        "abstract": "{In the context of direct imaging of exoplanets, coronagraphs are commonly proposed to reach the required very high contrast levels. However, wavefront aberrations induce speckles in their focal plane and limit their performance.} {An active correction of these wavefront aberrations using a deformable mirror upstream of the coronagraph is mandatory. These aberrations need to be calibrated and focal-plane wavefront-sensing techniques in the science channel are being developed. One of these, the self-coherent camera, of which we present the latest laboratory results.} {We present here an enhancement of the method: we directly minimized the complex amplitude of the speckle field in the focal plane. Laboratory tests using a four-quadrant phase-mask coronagraph and a 32x32 actuator deformable mirror were conducted in monochromatic light and in polychromatic light for different bandwidths.} {We obtain contrast levels in the focal plane in monochromatic light better than $3.10^{-8}$ (RMS) in the 5 -- 12 $\\lambda/D$ region for a correction of both phase and amplitude aberrations. In narrow bands (10 nm) the contrast level is $4.10^{-8}$ (RMS) in the same region.} {The contrast level is currently limited by the amplitude aberrations on the bench. We identified several improvements that can be implemented to enhance the performance of our optical bench in monochromatic as well as in polychromatic light.}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  Planets detected with indirect methods often orbit their star at a small semi-major axis, while direct imaging targets objects located at larger distances from their star. Therefore, the development of this technique might lead to the discovery of entirely new objects, following the detection of $\\beta$-pic b \\citep{Lagrange_Bpic09} and of the HR~8799 system \\citep{Marois08,Marois10}. Direct imaging performed on broad bandwidths will also allow surface and atmosphere spectroscopic analysis. However, the main challenge of direct imaging is the high contrast level required to separate the planet signal from the stellar light.  Most direct-imaging instruments use coronagraphs to reject the stellar light. But their performance is strongly limited by phase and amplitude aberrations in the incident wavefront, which produce stellar speckles in the image plane. To compensate for all these aberrations and reduce the speckle level, an active correction including a focal plane wavefront sensor paired with a deformable mirror (DM) is mandatory. Several techniques have been developed to retrieve phase and amplitude aberrations directly in the focal plane of a coronagraph, using either the application of known phases on the DM \\citep{BordeTraub06} or a specific instrumental design. Among these, the self-coherent camera \\citep[SCC, ][]{Baudoz06} uses the coherence between the focal plane speckles and the stellar light rejected by the coronagraph. This technique has already reached high contrast levels in simulations \\citep{Galicher10} and experimentally \\citep{Mazoyer13} in monochromatic light.  This paper presents the latest laboratory performance obtained with the SCC. An enhancement of the method is described in Section~\\ref{sec:theo_part}, and we present the current bench configuration in Section~\\ref{sec:setup}. Results are presented in Section~\\ref{sec:monochromatic} for monochromatic light and in Section~\\ref{sec:polyc} for polychromatic light. ",
        "conclusions": "\\label{sec:ccl}  We presented the latest improvements made to the self-coherent camera (SCC). We tested this technique on a laboratory bench and reached contrast levels (RMS) better than $3.10^{-8}$ in the 5 -- 12 $\\lambda/D$ zone. This means that we improved the contrast level by a factor of 13 with respect to the previous version of the SCC~\\citep{Mazoyer13}. Numerical simulations showed that the current limitation is imposed by uncontrolled amplitude defects. For wavelengths close to the optimal wavelength of the coronagraphic mask and for narrow bands, contrast levels measured in the laboratory are better than $4.10^{-8}$ in the 5 -- 12 $\\lambda/D$ zone. The current limitation is the chromaticity of our coronagraph. As expected, in broader bandwidths, a dark hole can be obtained in a smaller correction zone with degraded performance.  To enhanced the size and contrast level of this zone, we will study the multireference SCC, including experimental results, in a forthcoming paper. We also plan to use of the SCC with more achromatic coronagraphs, such as the multi-FQPM \\citep{Galicher11} or the dual-zone phase mask \\citep{Ndiaye12}. This is expected to improve the chromatic behavior and show that the SCC is compatible with several coronagraphic masks."
    },
    "1402/1402.7244_arXiv.txt": {
        "abstract": "{We seek to characterize giant-planet systems by their gravitational scattering properties. We do this to a given system by integrating it numerically along with a large number of hypothetical small bodies that are initially in eccentric habitable zone (HZ)-crossing orbits. Our analysis produces a single number, the escape rate, which represents the rate at which the small-body flux is perturbed away by the giant planets into orbits that no longer pose a threat to terrestrial planets inside the HZ. Obtaining the escape rate this way is similar to computing the largest Liapunov exponent as the exponential rate of divergence of two nearby orbits. For a terrestrial planet inside the HZ, the escape rate value quantifies the ``protective'' effect that the studied giant-planet system offers. Therefore, escape rates could provide information on whether certain giant-planet configurations produce a more desirable environment for life than the others. We present some computed escape rates on selected planetary systems, focusing on effects of varying the masses and semi-major axes of the giant planets. In the case of our Solar System we find rather surprisingly that Jupiter, in its current orbit, may provide a minimal amount of protection to the Earth. ",
        "introduction": "The fascinating dynamical variety of extrasolar planetary systems has been the motivation behind many recent numerical simulations. We have seen detailed studies concerning the orbital stability of observed multi-planet systems \\citep[e.g.,][]{BQ2004} and also stability analyses where fictitious terrestrial planets are integrated along with the observed giant-planet system \\citep[e.g.,][]{JUS2005,MT2003,ABC2004}. In a few cases (GJ 777A, 47 UMa, HD 4208, HD 72659 etc.) it is found that terrestrial planets could indeed survive inside the habitable zone (HZ) of the system in a million year time scale.  For those candidate systems that could harbor Earth-like planets, we can take our dynamical speculation one step further and ask: How intense would the small-body flux be in the habitable zone of the planetary system in question? After all, small-body impacts to the Earth have had a major influence on the evolution of terrestrial life. Although it is arguable whether a smaller or larger small-body flux would have been more beneficial for our evolution, it would still be a step forward if we could present something quantitative about the small-body fluxes in the extrasolar systems.  The initial conditions in the circumstellar disk, the process of planet formation, and subsequent dynamical evolution undoubtly sculpt an individual distribution of small bodies for each planetary system. A recent analysis of the debris disk around $\\tau$ Ceti \\citep{GWH2004} suggests that the disk has similar dimensions, but mass an order of magnitude greater than the Kuiper Belt, indicating a more intense flux of cometary bodies. Is seems that, again, our solar system is not necessarily to be taken as the prevalent specimen, and that the bombardment by small bodies may become an important variable when searching for habitable Earth-like planets. Direct observations of debris disks are currently limited to very nearby stars, and are unavailable for the majority of extrasolar planetary systems found by radial velocity measurements ({\\it http://www.obspm.fr/planets}). For these systems, models of planet formation could be used to produce a hypothetical distribution of small bodies, but this would lead us to undesirable complex simulations and heavy computational load. Hence, we accept the fact that their current population of small bodies is unknown, and suggest another kind of method of analysis.  In our approach we follow the evolution of a specific hypothetical population of small bodies which is common to all of the planetary systems under study. Our purpose is to isolate and characterize the small-body scattering properties of the planets, and to provide a technique to compare different planetary systems in this respect. If we can show that a certain configuration of planets is considerably more efficient in scattering small bodies than another one, we could argue that the actual small-body flux is less intense in that system, taken that the initial conditions in both systems are similar. \\citet{Wet1994} used this idea when he considered alternative giant-planet formation scenarios in our Solar System. Using {\\\"O}pik-Arnold calculations he followed the evolution of cometary test bodies and concluded that the absence of Jupiter would increase the cratering rate on Earth throughout its history by a factor of 100--1000. The dynamically dominant role of Jupiter in the Solar System was also shown by \\citet{Eve1968}, in the case of parabolic comets.  In this paper we will present computational tools to measure the gravitational scattering properties of a planetary system. First, in Sect.\\ \\ref{s:methods}, we will introduce a ``benchmark'' integration scenario with thousands of cometary small bodies to be integrated numerically along the planetary system under study. The small bodies are initially in orbits that cross the HZ of the system. We will also define a variable, the escape rate, that describes the strength at which the planetary system scatters the small bodies into orbits that no longer threaten the HZ. In Sect.\\ \\ref{s:results}, we will present some preliminary results obtained using the tools above. We have considered hypothetical giant-planet configurations based on the Solar System and especially evaluated the dynamical significance of Jupiter. Later, we will show some results on selected extrasolar planetary systems. In Sect.\\ \\ref{s:discussion} we will discuss the interpretation of the results. ",
        "conclusions": "\\label{s:discussion}  We have introduced a technique for measuring gravitational scattering efficiency in a planetary system. One simple dynamical quantity, the escape rate, describes the protection that a giant-planet system offers to a hypothetical terrestrial planet. Our results show that, at least with our current choice of parameters, the escape rate is a computationally consistent quantity and is a function of the orbital structure of the giant planets. The important question is: can we use escape rates to make conclusions that apply to real planetary systems or are they just a dynamical curiosity with no physical significance?  We believe that the very definition for the escape rate (Sect.\\ \\ref{s:escape_rate}) is rigorous and, when comparing planetary systems, the one with a higher escape rate would be more efficient in scattering small bodies away from the HZ. However, this argument applies only to the particular initial population of small bodies used in the escape rate scenario. Hence, it is the choice of the population that mostly determines what the escape rate really tells us.  There is an obvious conceptual difference that should be noted. In reality, the shape and intensity of the small-body flux in a planetary system is a function of time, whereas in the escape rate computation, the flux is intentionally kept constant. Therefore, each escape rate scenario represents a fixed moment in time. In order to simulate the small-body flux in a real planetary system, we should fix a time in the history (or in the future) of the system, and create an initial population that is a model for the HZ-crossing flux at that particular moment. Unfortunately, by doing so, we would inevitably lose the ability to compare the escape rates of totally different planetary systems. This kind of approach could still be valid and beneficial, if we concentrated to a one well known system (i.e., to our Solar System) and varied the parameters of the planets only by small amounts.  In the benchmark integration scenario presented in this paper (Sect.\\ \\ref{s:parameters}) we use the escape rate as a universal dynamical quantity which is not identified with real fluxes of small bodies. Our initial population is artificial, but for the sake of creditability, it should still be somehow ``typical'' among planetary systems.  Our distribution of particles is spherically symmetric, and in the orbital parameter space the particles cover the angular dimensions ($\\iota$, $\\Omega$, $\\tilde{\\omega}$, $L$) completely, but only a limited range in semi-major axis and eccentricity ($a$, $e$). This kind of distribution is not typical to the observed small-body populations or to the simulations explaining their evolution \\citep[see, e.g.,][ for Solar System comets]{DQT1987,DQT1988}. However, one should remember that the initial population for the escape rate is supposed to represent only the fraction of small bodies that threaten the HZ of the planetary system. The modelling of all the processes that bring small bodies into to HZ-crossing flux would be a cumbersome task and, hence, we feel justified to use our simplified initial population as a first order approximation.  As a part in improving the convergence and efficiency in escape rate computations, we chose the initial orbital distribution in a way that the apocenter distances ($Q$) were all well beyond the semi-major axis of the innermost giant planet. Obviously, a qualitatively different choice, e.g., one with the apocenter distances inside the orbits of the giants, could change our results (Sect.\\ \\ref{s:results}) significantly. This is something we wish to investigate in future papers. However, as an indication of partial robustness, the qualitative behaviour of our results did not change when we used an initial population with $Q=\\left[40,50\\right]$ and a narrower definition for the HZ and for the removal criterion by particle pericenter ($\\left[0.8,1.4\\right]$ and $1.5$ AU, respectively). Our choice for the initial population also ensures that the orbits of the particles are chaotic and none of the particles survive the whole integration time. If the population included numerous stable or resonant orbits, the regenerated particles might accumulate into safe areas in the phase space, and the escape rate might become biased.  There are open questions about the initial population, but our results with the benchmark scenario show that the escape rate has informational value, and it can be used in comparative analyses on planetary systems which was our goal in the first place. As such, the escape rate cannot predict the habitability of a hypothetical terrestrial planet in an observed extrasolar planetary system, but it could be used, among other tools, to make educated guesses about the matter.  The computed escape rate values for the Solar System giant planets confirm that Jupiter has the dominant role in scattering small bodies out of the HZ. The scattering efficiency in the Solar System would not be greatly disturbed if the other giant planets were removed. This is in line with the findings by \\citet{Wet1994}. On the other hand, removing Jupiter does not affect the escape rate considerably either since the other giants effectively take its role in scattering.  The minimum in the escape rate at the true semi-major axis of Jupiter is an interesting detail. It could be a coincidence, but perhaps there is some cosmogonical, or dynamical, explanation. Maybe Jupiter, the dominant mass, is in an optimal orbit to encourage the heavy bombardment of cometary bodies into the inner Solar System. This could have been an essential requirement for life and evolution on Earth.  If we compare the escape rate for the Solar System to the escape rates for \\object{47 UMa} and \\object{GJ 777A} we see that small-body scattering efficiency is more than an order of magnitude greater in the extrasolar systems. A natural explanation to this are the smaller semi-major axes and greater masses of the extrasolar giant planets. Comparison between \\object{47 UMa} and \\object{GJ 777A} is more interesting, because the higher escape rate of \\object{GJ 777A} is difficult to explain by intuition. By looking at the statistics produced by the integration, we saw that the outer planet in \\object{GJ 777A} was the most recent perturber for $92\\%$ of the scattered particles. Therefore, the large eccentricity of the outer planet is probably responsible for the higher escape rate value.  This is a preliminary paper where we have introduced a new concept, implemented a computational method, and demonstrated some basic applications. There are many possibilities for follow-up studies; an important one is to analyse further the effect of the initial population on the escape rate. Another task, which was beyond the CPU-time budget of this paper, would be to probe the parameter space of the Solar System giants more thoroughly. In addition to semi-major axes and planetary masses, one could include the eccentricities into the analysis. Generally, the increasing dimensionality of the parameter space will certainly become a problem, but with some kind of compromises, one could compute escape-rate maps which could reveal interesting dynamical structures."
    },
    "1402/1402.6597_arXiv.txt": {
        "abstract": "{For over a decade, the structure of the inner cavity in the transition disk of TW\\,Hydrae has been a subject of debate. Modeling the disk with data obtained at different wavelengths has led to a variety of proposed disk structures. Rather than being inconsistent, the individual models might point to the different faces of physical processes going on in disks, such as dust growth and planet formation.} {Our aim is to investigate the structure of the transition disk again and to find to what extent we can reconcile apparent model differences.} {A large set of high-angular-resolution data was collected from near-infrared to centimeter wavelengths. We investigated the existing disk models and established a new self-consistent radiative-transfer model. A genetic fitting algorithm was used to automatize the parameter fitting, and uncertainties were investigated in a Bayesian framework.} {Simple disk models with a vertical inner rim and a radially homogeneous dust composition from small to large grains cannot reproduce the combined data set. Two modifications are applied to this simple disk model: (1)~the inner rim is smoothed by exponentially decreasing the surface density in the inner $\\sim3\\,$AU, and (2)~the largest grains ($>100\\,\\mu$m) are concentrated towards the inner disk region. Both properties can be linked to fundamental processes that determine the evolution of protoplanetary disks: the shaping by a possible companion and the different regimes of dust-grain growth, respectively. } {The full interferometric data set from near-infrared to centimeter wavelengths requires a revision of existing models for the TW\\,Hya disk. We present a new model that incorporates the characteristic structures of previous models but deviates in two key aspects: it does not have a sharp edge at 4\\,AU, and the surface density of large grains differs from that of smaller grains. This is the first successful radiative-transfer-based model for a full set of interferometric data.} ",
        "introduction": "With the potential discovery of (proto)planetary bodies in disks around young stars \\citep{2012ApJ...745....5K,2013ApJ...766L...1Q}, observations started to provide the ultimate justification for calling the disks ``protoplanetary''. The conservation of angular momentum naturally explains why the collapsing cloud around stars evolves to a much more compact rotating disk of dust and gas. In contrast, the subsequent evolution of this protoplanetary disk towards a gas-free debris disk or a genuine planetary system is a much more complicated---and less understood---process (for an overview, see, e.g.,~\\citealt{2011ppcd.book..114H,2011ARA&A..49...67W}).  The current hypothesis on protoplanetary disk evolution includes a short dissipation phase, in which disk material disperses outwards. A substantial fraction of disks around young stars is indeed observed to have an inner gap or ``hole'', a population of so-called \\emph{transition disks}. At least four mechanisms have been proposed for the dissipation mechanism. First, \\emph{photoevaporation} might be responsible for physically clearing the inner-disk material, a process driven by energetic photons (e.g.,~\\citealt{2009ApJ...705.1237G,2010MNRAS.401.1415O}). A second mechanism is the \\emph{dynamical clearing} by a (sub-) stellar companion (e.g.,~\\citealt{1994ApJ...421..651A}). Third, \\emph{magneto-rotational instability} might take place in an ionized disk rim, and therefore drain out the inner disk (e.g.,~\\citealt{2007NatPh...3..604C}). Finally, the inner gap could actually only be an opacity gap caused by \\emph{grain growth}  (e.g.,~\\citealt{2005A&A...434..971D}), a process diminishing the emitting/absorbing surface of the dust and therefore creating (only) an \\emph{apparent} gap. Recent simulations \\citep{2011MNRAS.412...13O} suggest that a combination of several processes might be necessary to explain all cases.  The observational situation on transition-disk objects is complex. Recent mm-observations confirm low densities in mm-sized grains in the inner disk of a sample of transition disks \\citep{2011ApJ...732...42A}, but modeling of these data predicts much stronger depletion in $\\mu$m-sized grains than actually observed \\citep{2012ApJ...750..161D}. In the inner regions, several ``standard'' transition-disk sources seem to differ significantly. Based on scattered-light images and SED-modeling, \\citet{2012ApJ...752..143H} model the central cavity of HD\\,169142 as being almost completely empty, with only a very compact disk and a halo-like central region consisting of pure carbonaceous material. SED-modeling of the disk of LkCa 15 leads to similar geometries, but for this model, the compact component is predominantly filled with large silicate grains \\citep{2010A&A...512A..11M}. A third case is T Cha, for which \\citet{2011A&A...528L...6O} model near-IR interferometry and SED data with a tiny inner disk of small silicate and carbon grains. Similarly, HD\\,100546 is modeled as having a tiny inner disk, but the outer disk of this particular object is found to have a rounded rim \\citep{2013A&A...557A..68M}. Seemingly, for the handful of transition disks that have been modeled in detail, there is no consistent relation found on dust composition, grain sizes and inner-disk geometry.     Since it was identified to be a transition disk, the disk around TW Hydrae (TW\\,Hya) has been one of the principal objects for investigating the physical conditions in dispersing protoplanetary disks. Various studies reveal gas lines of different molecules, allowing us to constrain the physical properties of the gas in the disk (e.g.,~\\citealt{2010A&A...518L.125T,2011ApJ...735...90G,2011ApJ...736...13P,2012ApJ...748....6G,2012ApJ...757..129R}). The discovery of a handful of key molecules in the disk, such as gaseous H$_2$O \\citep{2011Sci...334..338H} and HD \\citep{2013Natur.493..644B}, allows pioneering work to be done in the study of early solar-system analogues. The latter work also led to an estimate of the total disk mass.   Several of the mentioned works start from a certain disk geometry for constraining model quantities. The geometry in its turn is derived from the dust emission profile, the most important diagnostic for constraining structural parameters of the disk. However, as we will indicate in Sect.~\\ref{sect:history}, a consensus on the disk structure is actually far from being reached. Different instruments operating at different wavelengths and resolutions probe different regions of the disk both in terms of radial distance to the central star and vertical distance above the disk midplane. In addition, they may be sensitive to very different parts of the dust (size) distribution. This can lead to large differences between models fit to different observations of the same object, even if each model adequately reproduces the data it was meant to explain. As a result, the full dust distribution in the disk, which forms the backbone for radiation balance and therefore the disk temperature structure, remains incompletely constrained.  We have collected an extended set of new and archival high angular resolution data for reinvestigating the TW\\,Hya disk. In Sect.~\\ref{sect:history} we give a short overview of disk-structure models that have been published for TW\\,Hya. Sect.~\\ref{sect:observations} introduces the data set and comments on the data reduction, while a first qualitative analysis of the data is given in Sect.~\\ref{sect:inspection}. We confront several previous disk models with the full data set in Sect.~\\ref{sect:reimplementing} and discuss potential shortcomings. This opens the way to Sect.~\\ref{sect:modeling}, where we present our detailed radiative-transfer analysis. In Sect.~\\ref{sect:discussion}, the results are being discussed, and Sect.~\\ref{sect:nearinfrared} goes in some more detail considering the analysis of the near-infrared data. Finally, Sect.~\\ref{sect:conclusion} summarizes the conclusions of this work. ",
        "conclusions": " \\begin{enumerate} \\item Reimplementing existing radiative-transfer models for the TW\\,Hya disk shows that the published models can indeed explain part of the combined data set, but fail to reproduce data at other wavelengths. This shows that a multiwavelength analysis is essential to fully characterize the disk. Also, this exercise already points to some model properties: the need for large ($\\gg\\mu$m) dust grains, the importance of full radiative transfer for explaining the disk, and the fact that the new mm-interferometry gives completely new constraints on the disk structure. \\item A simple disk structure with a vertical inner rim and a radially homogeneous composition (from small to large grains) cannot properly explain the data. From the MIDI data, we infer an inner region with a lower surface density. \\item As a first modification to the simple disk model, we introduce a rounded rim by exponentially reducing the surface density (with respect to a standard power law surface density) inwards of a specified radius. \\citet{2013A&A...557A..68M} have shown that this rim profile follows from hydrodynamical simulations of a low-mass companion opening a disk gap. Our MIDI data agree with the expected mid-infrared emission of such a rim. Although linking the gap in the TW\\,Hya disk with the presence of a companion is tentative, other mechanisms might lead to similar rim geometries. \\item The mm-emission as observed with VLA originates from a spatially more compact region than the sub-mm emission observed with SMA. We interpret this as evidence for a more compact distribution of the largest dust grains, something that is also seen in other disks (e.g.,~\\citealt{2012ApJ...760L..17P}). We model the $>100$-$\\mu$m grains with the same surface density profile as the smaller grains, but with a steeper surface density power. This allows us to reproduce both the SMA- and VLA-visibilities, to a reasonable extent. We note that for TW\\,Hya, the VLA data allow us to trace the imprint of dust emission to wavelengths as long as 4 cm, in line with earlier suggestions by \\citet{2005ApJ...626L.109W}. \\item Our final disk model (Fig.~\\ref{fig:final_drawing}) has an inner radius of $0.3-0.5\\,$AU, an outer radius of $\\sim60\\,$AU, and a maximum in the surface density around $2.5\\,$AU, where the surface density smoothly decreases inwards. \\item For the radiative-transfer analysis, we do not focus on the near-infrared data, since these depend predominantly on the scattering properties of the dust. Instead, we compare the visibilities predicted by the best-fit model with the actual data. This comparison shows the importance of taking the full scattering properties of the dust into account. The predicted visibilities reproduce the H-band data reasonably well; for the L'-band data, a small population of possibly fluffy grains might be needed in the disk atmosphere for predicting the correct visibilities. A population of small grains at large distances from the star, as observed in scattered light images, can easily be incorporated in the final disk model. \\end{enumerate}"
    },
    "1402/1402.4889_arXiv.txt": {
        "abstract": "High-resolution astronomical imaging at sub-GHz radio frequencies has been available for more than 15 years, with the VLA at 74 and 330 MHz, and the GMRT at 150, 240, 330 and 610 MHz. Recent developments include wide-bandwidth upgrades for VLA and GMRT, and commissioning of the aperture-array-based, multi-beam telescope LOFAR. A common feature of these telescopes is the necessity to deconvolve the very many detectable sources within their wide fields-of-view and beyond. This is complicated by gain variations in the radio signal path that depend on viewing direction. One such example is phase errors due to the ionosphere.  Here I discuss the inner workings of SPAM, a set of AIPS-based data reduction scripts in Python that includes direction-dependent calibration and imaging. Since its first version in 2008, SPAM has been applied to many GMRT data sets at various frequencies. Many valuable lessons were learned, and translated into various SPAM software modifications. Nowadays, semi-automated SPAM data reduction recipes can be applied to almost any GMRT data set, yielding good quality continuum images comparable with (or often better than) hand-reduced results. SPAM is currently being migrated from AIPS to CASA with an extension to handle wide bandwidths. This is aimed at providing users of the VLA low-band system and the upcoming wide-bandwidth GMRT with the necessary data reduction tools. ",
        "introduction": "\\label{sec:intro}  Low-frequency radio telescopes have relatively large fields-of-view, because the telescope primary beam size scales with wavelength. Current high-resolution (sub-arcminute) radio interferometers that operate at sub-GHz frequencies, like the Giant Metrewave Radio Telescope (GMRT), the LOw-Frequency ARray (LOFAR), or the Very Large Array (VLA), have field diameters measured in (sometimes tens of) degrees. Typical challenges of handling high-resolution, low-frequency continuum observations are (i) imaging and deconvolving many sources inside (and outside!) the large primary beam area, (ii) direction-dependent (DD) visibility amplitude and phase variations across the field-of-view due to antenna beam patterns, pointing errors, ionosphere, etc., and the abundance of Radio Frequency Interference (RFI).  Processing data coming from these telescopes can be daunting, since all the challenges mentioned above (and more) needs to be dealt with at the same time. In data reduction, the \\emph{art} is to determine at any given time which effect is dominant in limiting the image quality. Generally, a few \\emph{big effects} stand out and are generally easy to identify and fix. But subsequenty, many \\emph{smaller effect} will present themelves, which may be more difficult to recognize, and harder to identify and disentangle. A effective way to mitigate the smaller effects is to perform iterations of (DD) calibration, imaging and flagging, in each step refining the quality of both the visibility data and the reconstructed sky model. ",
        "conclusions": "\\label{sec:summary}  SPAM is an AIPS-based Python package that provides semi-automated data reduction scripts for all sub-GHz frequencies at GMRT. Apart from well-tested standard data reduction steps, SPAM includes direction-dependent (ionospheric) calibration and imaging, and alternative RFI- and bad data mitigation methods. With SPAM, GMRT data reductions are highly efficient, highly reproducible, and give good quality results. SPAM is available for everyone to use\\footnote{\\url{https://safe.nrao.edu/wiki/bin/view/Main/HuibIntemaSpam}}.  SPAM for CASA, currently under development, will also provide wide-bandwidth ionospheric calibration, direction-dependent ionospheric Faraday rotation calibration, and will naturally benefit from the advanced modes (and continuous improvements) of the CASA imager \\citep[e.g., see][and references therein]{2013ApJ...770...91B}. %"
    },
    "1402/1402.4137_arXiv.txt": {
        "abstract": "The nature of the Type Ia supernovae (SNIa) progenitors remains still uncertain. This is a major issue for galaxy evolution models since both chemical and energetic feedback play a major role in the gas dynamics, star formation and therefore in the overall stellar evolution. The progenitor models for the SNIa available in the literature propose different distributions for regulating the explosion times of these events. These functions are known as the Delay Time Distributions (DTDs). This work is the first one in a series of papers aiming at studying five different DTDs for SNIa. Here, we implement and analyse the Single Degenerate scenario (SD) in galaxies dominated by a rapid quenching of the star formation, displaying the majority of the stars concentrated in the bulge component.  We find a good fit to both the present observed SNIa rates in spheroidal dominated galaxies, and to the [O/Fe] ratios shown by the bulge of the Milky Way.  Additionally, the SD scenario is found to reproduce a correlation between the specific SNIa rate and the specific star formation rate, which closely resembles the observational trend, at variance with previous works. Our results suggest that SNIa observations in galaxies with very low and very high specific star formation rates can help to impose more stringent constraints on the DTDs and therefore on SNIa progenitors. ",
        "introduction": "\\label{Sec:Intro}  Galaxy formation represents a multi-scale highly non-linear process. The modeling of the observed galaxy populations with a self-consistent model -- from the initial conditions left behind the Big Bang--, requires knowledge of the co-evolution of all the components of the Universe. According to the $\\Lambda$-Cold Dark Matter ($\\Lambda$CDM) paradigm, baryons condensate onto the dark matter haloes, which constitute the sites of galaxy formation. The theories of galaxy formation aim at understanding at the same time the large scale growth of the structure and small-scales processes such as the star formation from molecular clouds, for example. This is a challenging task which can be tackled by using semi-analytical or fully cosmological models. Both, resort to recipes or sub-grid modeling to include complex physical processes which cannot be numerically resolved. Hence, is of utmost importance to confront the model results with observations to learn about galaxy formation and to test the validity of the adopted hypotheses.   Galaxy chemical evolution models provides us with a powerful tool to understand the way in which galaxies formed and evolved \\citep{Tinsley79,MatteucciGreggio86,Matteucci94}.  Since the information given by chemical abundances is the result of the nucleosynthesis production of the stellar populations and large scale physics involved in the galaxy formation, chemical patterns can provide stringent constraints for galaxy formation models.  \\begin{sloppypar} The treatment of chemical evolution with the inclusion of different models for Type Ia supernovae (SNeIa) opened the route to study galaxy formation by probing links between chemical patterns and relevant physical processes.  These efforts are usually attempted with a variety of numerical techniques, such as multi-zone chemical evolution models \\citep[e.g.][]{MatteucciGreggio86,Isern93,Cris01,PM04}, semi-analytic models of galaxy formation ~\\citep[e.g.][]{CaluraMenci09,Arrigoni10,Jimenez11,Yates13, DeLucia14, Tuna15}, and cosmological hydrodynamical simulations ~\\citep[e.g.][to name a few]{Raiteri96, Carraro98,Mosconi01,Tornatore04,Nakasato03, Kobayashi04,SC05,Nagashima05,Wiersma09, Aumer13, Few14}.  Nowadays, sophisticated computational tools and detailed observations of local and high redshift galaxies enable us to improve galaxy formation models and develop more realistic schemes for the baryonic astrophysics. \\end{sloppypar}  One of the major astrophysical problems is the uncertain identity of the SNIa progenitors. This is a matter of concern not only in galaxy formation theory but also in modern cosmology. The remarkable similarity shown by the light curves of the SNIa have made them excellent cosmological distance indicators. Hence, they are used to investigate the properties of dark energy, and to test parameters of the cosmological model and the acceleration epoch of the Universe \\citep{Perlmutter99}. However, given the still unknown nature of the progenitor systems of SNIa, systematic errors in the deduced distances based on calibrations of nearby SNeIa, could have an impact in the cosmological parameter estimation \\citep{PanDulce12, Rest13}.  It is well known that core-collapse supernovae (SNeII) are produced by massive short-lived stars (M$>8$\\Msun). The nucleosynthesis production of these events feeds the interstellar medium (ISM) with energy and, mainly, with the so-called $\\alpha$-elements -- O, Ne, Mg, Si, S and Ca --. On the other hand, the SNeIa are the main contributor of the Fe in the Universe \\citep{Greggio83, MatteucciGreggio86, Cappellaro97}. This element is usually considered to be a tracer of the metallicity in stars.  Furthermore, the production and distribution of the chemical elements in the ISM affects the cooling rates of the gas \\citep{ShuterlandDopita93}, the star formation processes, the subsequent stellar evolution \\citep{Pietrinferni06}, and the production of high-redshift dust \\citep{Maiolino04,Bianchi07}. Main properties such as the luminosity function and the mass distribution of galaxies might consequently be affected \\citep{WhiteFrenk91,SC05}. By studying the enhancement of the $\\alpha$-elements relative to Fe, we expect to learn about the Initial Mass Function (IMF) of the stellar populations and, very importantly, about the time-scales in which SNeIa become relevant. The time-delay between core-collapse SNe and SNeIa injection of chemical elements in the ISM, creates important patterns which can be used as clocks to tag particular events in the formation histories of galaxies.  Given a starburst, SNIa events occur following a distribution of explosion times which is known as the Delay Time Distribution (DTD).  Different scenarios for SNIa progenitors produce different DTDs which, in turn, create particular chemical patterns in the stellar populations and their host galaxies.    It has long been known \\citep{HoyleFowler60} that the explosion mechanisms, the chemical and energetic composition of the remnants and the light curves of SNeIa (with absence of He and H) involve, at the most fundamental level, the combustion of a degenerate stellar core. Specifically, a White Dwarf (WD) star of Carbon (C) and Oxygen (O) is lead away from the equilibrium, following mass accretion and successive explosion.  Among the two most popular scenarios for explaining SNIa based on a thermonuclear explosion of a C-O WD, there is the Single Degenerate (SD) scenario -- a WD exceeding the Chandrasekhar mass through accretion from a non-degenerate companion star-- where the mass accretion can assume many configurations. The secondary star is proposed to be a Main Sequence star, a subgiant star, a helium star or red giant star \\citep{WhelanIben73,IbenTutukov84,VandenHeu92,Hachazo96,Hachazo99,Hansolo04,Geier13}. In addition to the variety of possible secondary stars, some models include other effects such as the metallicity dependence proposed by \\cite{Kobayashi98}. The authors found that when the metallicity is lower than [Fe/H]$< -1$, the winds developed by the primary star (WD) accreting mass through the Roche lobe are too weak and the explosion cannot occur.  The second most popular scenario involving a thermonuclear explosion of a C-O WD is the Double Degenerate (DD) scenario -- two WDs that loose angular momentum and energy by emitting gravitational waves and eventually merge. If the exceed the Chandrasekhar mass they ignite as SNIa \\citep{IbenTutukov84}.  Other suggested theoretical scenarios include the ``Collisional Scenario'' where the head-on collisions of two WDs occurs instead of the spiralling due to loss of gravitational wave radiation. These collisions of two WDs of sufficiently large masses are predicted in dense environments such as globular clusters, and could explain supernovae occurring in the nuclei of galaxies \\citep{LorenIsernyelP10}. On the other hand, the ``Core-Degenerate (CD) scenario'', where a WD merges with the core of an asymptotic giant branch (AGB) star and forms a rapidly rotating WD. This new configuration has a mass close to and above the critical mass for explosion, and was used recently as the best scenario to explain the observed properties of SN 2011fe \\citep{Althaus13}. Additionally, the ``Double Detonation'' mechanism propose a sub-Chandrasekhar WD accumulating a layer of helium-rich material on the surface sufficiently massive and degenerate to cause a detonation \\citep{Shen13}. Finally, another type of Core-Degenerate scenario considered is the ``super Chandrasekhar scenario''. These models propose a Chandrasekhar or a super-Chandrasekhar mass WD formed in the planetary nebula phase (or at the end of the common envelope phase), from a merger of a WD companion with the hot core of a massive asymptotic giant branch (AGB) star \\citep{Tenebroso13}.  However, a common characteristic of all these scenarios is that similar nucleosynthesis processes occur during the explosion creating mainly Fe. For a recent review on the SNIa progenitors see \\cite{Maoz13}.   Lately, many empirical DTDs have been proposed. The Bimodal model by \\citet{Mannucci06} considers a DTD with two populations of progenitors of SNIa, one dominated by the ``prompt'' component that explodes within 100 Myrs after the formation of their progenitors, and the ``tardy'' component exploding on a wide period of time extending up to 10 Gyr .  The shape of the bimodal distribution is given by the sum of two functions: a Gaussian centered at $5\\times10^{7}$yrs, and an exponentially declining plateau function. Moreover, similar power law DTDs ($\\sim$t$^{-1}$) have been recovered from different observational surveys. For instance, \\cite{Totani08} using field elliptical galaxies proposed such a power law. \\cite{Maoz12} also suggested a very similar DTD based on a sample of 132 SNeIa from the Sloan Digital Sky Survey II. Finally, \\cite{Pritchet08} presented a power law DTD ($\\sim$t$^{-0.5}$) from the SuperNova Legacy Survey.  The DTD represents a powerful tool, not only for testing models and helping to constraint the progenitor SNIa scenarios, but also for obtaining an accurate description of the chemical and energetic feedback from SNIa. The aim of our work is to test which is the best scenario for SNIa in the framework of cosmological simulations.  In particular, we will include and analyse in detail, for the first time to our knowledge, five different DTDs for SNIa in a smooth particle hydrodynamical code.  Our results will be presented in a series of three papers.  In this first paper we will discuss in detail the SD model and compare our results, obtained for a bulge-type galaxy, with the chemical abundances and SNIa rates for this kind of galaxy. In particular, we will compare our results with the abundance data for the Milky Way bulge.  We study global properties such as the correlation found by \\citet{Sullivan06} between the specific star formation rate (SSFR) and the specific SNIa rate (SSNIaR).  The reason for starting with this particular scenario is that it has been suggested as one the best models to reproduce the chemical properties of galaxies \\citep{Matteucci06,MatteucciFelice09}.   \\begin{figure} \\begin{center} \\includegraphics[width=0.55\\textwidth]{TablaSD.eps} \\includegraphics[width=0.55\\textwidth]{figHist.eps} \\caption{ ({\\it Upper panel}) The full DTD for the SD scenario given by \\citet{MR01} and the DTD parametrized with 20 bins to reduce the computational costs (for more details, see the text). ({\\it Lower panel}) Distribution of [Fe/H] ratios for the stars in the simulation, presented as an example of the relevant quantities of this study not being affected by the adopted number of bins used in the parametrization of the SD.} \\label{figConv} \\end{center} \\end{figure}    This paper is organized as follows: Section \\ref{Sec:DTD} describes the analytic DTD for the SD scenario. Section \\ref{Sec:code} briefly describes the main aspects of the numerical code and presents the initial conditions. In the following Subsections we present the implementation of DTD in the code. In Section \\ref{Sec:calibre} we show the calibration of the model with the observables. In Section \\ref{Sec:correlacion} we study the observed correlation found by \\citet{Sullivan06}, between the specific SNIa rates and the specific SFR for the SD scenario.  The main conclusions are summarized in Section \\ref{Sec:Conclusions}. ",
        "conclusions": "\\label{Sec:Conclusions}   We present results from simulated galaxies performed with a version of {\\small Tree-PM SPH-GADGET-3} \\citep{Springel05, SC06}, to study the impact of the SNIa feedback by using different DTDs. In this paper, we explore the SD model of MR01 in detail.  We choose the SD scenario as a reference to discuss the implementation, the main observables and correlations which we aim to reproduce.  The implementation of the SD scenario involves the calibration of a free parameter A which represents the fraction of binary systems in one stellar generation giving rise to SNIa events. Chemical evolution models generally fix A according to the SFR and the present day SNIa rates of the galaxy under study. However in a cosmological simulations A acts at a particle basis. Thus, it is assumed to be the same for all single stellar populations, while the final SNIa rates or any relation between the SNIa rates and the galaxy properties -- or the chemical abundance patterns -- should come out as a prediction of the simulations.  Therefore, it is necessary to explore a range of suitable A values and their impact on the properties of galaxies, by using simple initial conditions. In a forthcoming paper, we explore other progenitor scenarios. The final goal is to run cosmological simulation and analyse the chemical properties of galaxies of different masses and assembly histories.  Adequately reproducing the SNIa feedback as well as their effect on the chemical enrichment is very important, as new Galactic surveys will start yielding high precision measurements of chemical abundances.    Our main results can be summarized as follows:  \\begin{itemize} \\item We find that the SFR responds linearly to the number of SNeIa in the SD scenario. The SFR is found to decline faster with increasing values of A. But the differences are small and so the final stellar mass formed is very similar.  There is a substantial difference between the original model of \\citet{SC06} and the SD of \\citet{MR01} implemented and analysed here. This is related to the adopted lifetimes for the secondary masses. In the original model, the SNIa production was shifted to later times, compared to the SD scenario, beginning at $\\sim1\\times10^{8}$yrs. This excludes the so-called prompt SNeIa, that are indeed observed. For a similar A, the original model shows a more extended SFR. However, both of them predict similar final fractions of new stellar mass to total baryonic mass.\\\\  \\item The SFR in the simulations is dominated by a strong starburst consuming the cold gas within $\\sim1$ Gyr and forming stars that end up concentrated within the bulge.  Hence, the SNIa rates are comparable to observed ones for spheroidal galaxies. This comparison with observed SNIa rates of \\cite{Li11} shows that the best fits are obtained for SD-1 and SD-4 models, adopting A$\\sim 0.0015$ and A$\\sim 0.00075$, respectively.\\\\  \\item The [O/Fe] ratios predicted for the stars in the bulge of the simulated galaxies are compared to data from the Galactic Bulge (although this comparison is just indicative). We find that the best agreement with observations is again provided by the SD-1 and SD-4 models. These models predict a long plateau for the [O/Fe] ratio and a knee occurring at high [Fe/H], as observations suggest. We note that this is the first time that such an excellent agreement is found in galaxy simulations.\\\\  \\item At variance with previous claims, we find that the SD scenario (SD-1 and SD-4 models) reproduce the observed correlation between the SSNIaR and the SSFR, found by \\citet{Sullivan06}, if we estimate these quantities at different evolutionary times. This correlation comes out naturally from the simulations. Moreover, two features in the correlations are shown which cannot be confronted with current observations but could be interesting to explore in the future:  \\begin{itemize} \\item 1) For high SSFR, as the SSNIaR starts to appear, the SSNIaR-SSFR anti-correlates before it turns into the observed correlation.  This is caused by the large initial number of SNe Ia in the assumed DTD. The turnover occurs when the maximum in the DTD is reached and hence, it could be interesting to explore if this turn-over could be confirmed with observations coming from galaxies dominated by a very recent starburst.  \\item 2) For very low SSFR, there is still a residual SSNIaR because the majority of these events are delayed in time, thus the correlation disappears. The SSFR at which this occurs could be related to the shape of the DTD.  \\end{itemize} \\end{itemize}    In a forthcoming paper, we will discuss other DTDs implemented in the SPH simulations: the DD scenario by \\citet{Greggio05} and empirically motivated DTDs such as the Bimodal scenario by \\citet{Mannucci06}, the power laws reported by \\citet{Maoz12} and \\citet{Pritchet08}.  We will explore the predicted SNIa rates and chemical abundances, as well as the global correlations which can contribute to the understanding of SNeIa progenitors."
    },
    "1402/1402.2846_arXiv.txt": {
        "abstract": "We estimate dark matter density for the Universe with a reheating temperature smaller than the mass of dark matter, assuming dark matter to be a weakly interacting massive particle. During the reheating process, an inflaton decays and releases high-energy particles, which are scattered inelastically by the thermal plasma and emit many particles. Dark matters are produced through these inelastic scattering processes and pair creation processes by high-energy particles. We properly take account of the Landau-Pomeranchuk-Migdal effect on inelastic processes and show that the resultant energy density of dark matter is much larger than that estimated in the literature and can be consistent with that observed when the mass of dark matter is larger than $O(100)$ GeV. ",
        "introduction": "Introduction}  A weakly interacting massive particle (WIMP) is one of the most attractive candidates for dark matter (DM), motivated by new physics at a $\\TeV$ scale, including supersymmetric (SUSY) theories. DM is produced thermally in the early Universe, and its abundance can be consistent with that observed, if the reheating temperature of the Universe is sufficiently larger than its freeze-out temperature. The coincidence of the observed dark matter density and the relic density determined by the weak interaction scale is referred to as the WIMP miracle. Since this scenario requires the mass of the WIMP at the electroweak or $\\TeV$ scale, there are rich implications for near-future experiments, including direct and indirect detection experiments of DM and particle collider experiments.   However, when we look at each model motivated by particle physics, it is nontrivial to obtain a correct mass spectrum that can account for the abundance of DM. For example, in the constrained minimal SUSY standard model, the lightest SUSY particle (LSP) is bino-like. Non-observation of SUSY particles and the discovery of the 126~GeV Higgs boson~\\cite{Aad:2012tfa,Chatrchyan:2012ufa} by the LHC experiment indicate that SUSY particles are heavy, which leads to overproduction of the bino-like LSP in the early Universe.% \\footnote{ This problem can be avoided by co-annihilation~\\cite{Griest:1990kh,Gondolo:1990dk} with the stau~\\cite{Ellis:1998kh}, but only with fine-tuning. For recent discussion on SUSY models with a correct DM abundance, see Refs.~\\cite{PGM,Feng:2011aa,Okada:2012gf,Feng:2012rn,ArkaniHamed:2012gw,Yanagida:2013ah}, for example. } Taking this situation seriously, we reconsider the assumption of high reheating temperature.   When we consider inflationary models, a scenario with a low reheating temperature is naturally realized as follows. The inflaton is required to have a very flat potential, which suggests some symmetry to control its potential. The symmetry naturally suppresses interactions of the inflaton, which in turn leads to a low reheating temperature of the Universe. For example, if the mass of the inflaton is of the order of $10^{11}$ GeV and it decays through a dimension six  Planck-suppressed operator, the reheating temperature of the Universe is less than about $O(1)$ GeV, which is smaller than typical freeze-out temperatures of WIMPs.   In a scenario with a low reheating temperature, the thermal abundance of DM is much less than a scenario with high reheating temperature, mainly because the energy density of the thermal plasma is a subdominant component of that of the Universe at the time of DM freeze out~\\cite{Chung:1998rq, Giudice:2000ex}. In other words, the thermal abundance of DM is diluted by the entropy production from the inflaton decay.   However, the entropy production itself provides DM~\\cite{Moroi:1994rs, Kawasaki:1995cy, Moroi:1999zb, Allahverdi:2002nb, Gelmini:2006pw, Kurata:2012nf}. In Ref.~\\cite{Kurata:2012nf}, they have indicated that DM is produced in a shower from the decay of the inflaton and have calculated the resultant DM energy density using generalized Dokshitzer-Gribov-Lipatov-Altarelli-Parisi (DGLAP) equations~\\cite{Sarkar:2001se, Barbot:2002ep} in a certain SUSY model. They have found that DM is produced efficiently through this process when the inflaton decays into particles carrying a non-zero ${\\rm SU}(3)_{\\rm c}$ charge. The DM abundance also depends on the mass of inflaton, and the number of DM produced per one inflaton decay is typically $O(100)$ for the inflaton mass of $O(10^{12})$ GeV. In addition, in Ref.~\\cite{Allahverdi:2002nb}, it has been pointed out that DM is produced by inelastic scatterings between the thermal plasma and high-energy particles produced by the inflaton decay.   Therefore, in a scenario with low reheating temperature, the total amount of DM is given by the sum of the following contributions: (suppressed) thermal production, production through a cascade shower from the inflaton decay, and production through inelastic scatterings between high-energy particles and the thermal plasma. The last contribution is closely related to thermalization processes in the era of reheating, which has to be investigated in detail.   When reheating temperature is low enough, the typical momentum of particles produced by the decay of inflaton, which is roughly given by the mass of the inflaton, is much larger than the (would-be) temperature of background plasma. In this case, the thermalization process is completed by splitting processes through which the number of high-energy particles drastically increases~\\cite{Kurkela:2011ti,Harigaya:2013vwa} (see also Refs.~\\cite{Davidson:2000er, Arnold:2002zm,Jaikumar:2002iq}). The rate of splitting processes is suppressed by the Landau-Pomeranchuk-Migdal (LPM) effect, which is a destructive interference effect between emission processes~\\cite{Landau-Pomeranchuk, Migdal, Gyulassy:1993hr, Arnold:2001ba, Arnold:2001ms, Arnold:2002ja, Besak:2010fb}. The increases in the number of high-energy particles as well as the LPM effect should be taken into account in the estimation of the DM abundance produced through inelastic processes.   In this paper, we calculate the abundance of DM produced from inelastic scatterings between high-energy particles and the thermal plasma with careful consideration on the thermalization process as mentioned above. The resultant DM abundance is independent of the mass of the inflaton as long as the mass is sufficiently large, and depends mainly on the mass scale of the DM sector and reheating temperature. We find that this is the dominant contribution to the amount of DM in a scenario with a low reheating temperature when the mass of the inflaton is sufficiently large. We should emphasize that this mechanism to produce DM is highly model independent. Even if the decay channel of the inflaton into DM is absent in particular, DM is produced through inelastic scatterings. In addition, this scenario can also account for the abundance of DM with mass of $O(1)$ PeV, which is larger than unitarity bound of a few hundred TeV~\\cite{Griest:1989wd}. Such heavy DM might account for the recent observation of high-energy cosmic-ray neutrinos by the IceCube experiment~\\cite{Aartsen:2013jdh, Feldstein:2013kka, Esmaili:2013gha, Bai:2013nga}.   This paper is organized as follows. In the next section, we explain how high-energy particles lose their energy in the thermal plasma taking account of the LPM effect. We also describe a thermal history of the Universe in our setup. In Sec.~\\ref{production}, we briefly review previous works for the thermal and non-thermal production of DM and improve the calculation of the DM abundance from the inflaton decay taking account of the LPM effect. Then we discuss the relation between our scenario and other topics, such as the free-streaming velocity of DM, the Affleck-Dine baryogenesis, and SUSY theories. Section~\\ref{conclusion} is devoted to the conclusion. ",
        "conclusions": "conclusions}   We have considered WIMP DM in a scenario with low reheating temperature. Although there are several mechanisms to produce DM in this scenario including thermal production and production in a shower from the decay of the inflaton, the DM production by inelastic scatterings between high-energy particles and the thermal plasma gives the dominant contribution when the mass of the inflaton is sufficiently large. We have found that the abundance of DM depends mainly on the mass scale of the DM sector and the temperature of reheating, but not on the mass of the inflaton as long as the mass is sufficiently large. We have also found that the observed DM abundance can be accounted for when DM is heavier than $O(100)$ GeV. This conclusion is highly independent of the branching ratio of the inflaton as long as it decays into Standard Model particles since high-energy particles split into a lot of particles throughout inelastic scattering processes and the information of the initial condition is lost.   The above scenario is related to some cosmological topics. For example, the recent observation of high-energy cosmic-ray neutrinos by the IceCube experiment may indicate that DM is a long-lived particle with a mass of $O(1)$ PeV. The above scenario for non-thermal production of DM can also account for the abundance of even such heavy DM. In addition, since DM is produced non-thermally after the time of DM decoupling, it might have a cosmologically relevant free-streaming velocity. If interactions between DM and the thermal plasma are irrelevant after the DM production, the present-day free-streaming velocity of DM is non-zero and would be detected by future observations of the redshifted 21-cm line."
    },
    "1402/1402.5182_arXiv.txt": {
        "abstract": "The C--to--O ratio is a crucial determinant of the chemical properties of planets. The recent observation of WASP 12b, a giant planet with a C/O value larger than that estimated for its host star, poses a conundrum for understanding the origin of this elemental ratio in any given planetary system. In this paper, we propose a mechanism for enhancing the value of C/O in the disk through the transport and distribution of volatiles. We construct a model that computes the abundances of major C and O bearing volatiles under the influence of gas drag, sublimation, vapor diffusion, condensation and coagulation in a multi--iceline 1+1D protoplanetary disk. We find a gradual depletion in water and carbon monoxide vapors inside the water's iceline with carbon monoxide depleting slower than water. This effect increases the gaseous C/O and decreases the C/H ratio in this region to values similar to those found in WASP 12b's day side atmosphere. Giant planets whose envelopes were accreted inside the water's iceline should then display C/O values larger than those of their parent stars, making them members of the class of so-called ``carbon-rich planets''. ",
        "introduction": "The C/O ratio is a key parameter for the chemical composition and evolution of giant planets atmospheres since it controls the relative abundances of C-- and O--bearing species. Studies suggest that at equilibrium, as the C/O ratio increases ($\\geq$0.8) in the gas phase, all the available O goes into organics, CO, CO$_2$ and CH$_3$OH, so that the gas phase becomes H$_2$O-free and the remaining C is in the form of CH$_4$ \\citep{2010ApJ...709.1396F,2011ApJ...743..191M}. The CO/CH$_4$/CO$_2$ ratios in planetary atmospheres are also affected by the possible existence of non equilibrium chemistry effects due to dynamical mixing and photochemistry \\citep{2012A&A...546A..43V}. The C/O ratio is also crucial for understanding the chemical evolution of protoplanetary disks \\citep{1999JGR...10419003C,1993prpl.conf.1005P}.  Based on comparisons between a grid of model predictions and six narrowband photometry points for the dayside atmosphere of the Hot Jupiter (HJ) WASP 12b, \\cite{2011Natur.469...64M} found a best fitting model of C/O$\\geq$1 and C/H = 2$\\times$10$^{-5}$--1$\\times$10$^{-3}$, to compare with the stellar host values of C/O $\\sim$ 0.56 and C/H $\\sim$ 6$\\times$10$^{-4}$. These authors concluded that WASP 12b probably has a superstellar C/O but with a substellar C/H ratio. Because giant planets are expected to have stellar or superstellar metallicities \\citep{2005astro.ph..4214K}, new formation models consistent with the properties of WASP 12b have been proposed. For instance, \\cite{2011ApJ...743L..16O} considered the H$_2$O, CO and CO$_2$ snowlines in a protoplanetary disk model, and found an increase of the gaseous C/O ratio (compared to the C/O ratio of the stellar host) in the disk regions boxed by these snowlines, even reaching $\\sim$1 between the CO and CO$_2$ snowlines. They proposed that WASP 12b might have accreted its gaseous envelope in this area but they neglected the dynamical and thermodynamic processes {like gas drag, diffusion and sublimation} that may affect the abundances of volatiles in protoplanetary disks. On the other hand, \\cite{2011ApJ...743..191M} have shown that WASP 12b's atmospheric C/O is inconsistent with an atmosphere of the host star's composition, even if enriched in heavy elements by icy planetesimals dissolved in the planet's envelope, and suggest that it may have been formed in an O--depleted part of the protoplanetary disk. %  Here, we present a fully dynamical and multi--snowline model describing the transport and distribution of C and O bearing volatiles in protoplanetary disks, which is inspired from the original idea of \\cite{1998Icar..135..537C}. The model is used to compute the C/O ratio as a function of the radial distance in order to constrain the possible formation locations of carbon-rich HJs\\footnote{In this manuscript, carbon-rich refers to high C/O ratio, regardless of the C/H value.}. In Sec. 2 we present the different components of our model. Our results are presented in Section 3. In section 4 we discuss the implications and predictions of our model for the formation of carbon-rich HJs, and in particular for WASP 12b. ",
        "conclusions": ""
    },
    "1402/1402.0257.txt": {
        "abstract": "The Luminous Blue Variable stars exhibit behavior ranging from light curve `microvariations' on timescales of tens of days, to `outbursts' accompanied by mass loss of $\\sim10^{-3 }$ \\msun\\ occurring decades apart, to `giant eruptions' such as seen in Eta Carinae ejecting one or more solar masses and recurring on timescales of centuries.  Here we review the work of the Los Alamos group since 1993 to investigate pulsations and instabilities in massive stars using linear pulsation models and non-linear hydrodynamic models.  The models predict pulsational variability that may be associated with the microvariations.  Using a nonlinear pulsation hydrodynamics code with a time-dependent convection treatment, we show that, in some circumstances, the Eddington limit is exceeded periodically in the pulsation driving region of the stellar envelope, accelerating the outer layers, and perhaps initiating mass loss or LBV outbursts.  We discuss how pulsations and mass loss may be responsible for the location of the Humphreys-Davidson Limit in the H-R diagram.  The `giant eruptions', however, must involve much deeper regions in the stellar core to cause such large amounts of mass to be ejected.  We review and suggest some possible explanations, including mixing from gravity modes, secular instabilities, the epsilon mechanism, or the SASI instability as proposed for Type II supernovae.  We outline future work and required stellar modeling capabilities to investigate these possibilities. ",
        "introduction": "\\label{intro}  The Luminous Blue Variables (LBVs) are a class of high-mass stars that undergo eruptive events accompanied by mass loss.  Many have been discovered in our galaxy, but many are also found in the Large Magellanic Cloud \\citep[e.g.][]{wolf00}.  The LBVs show three different types of variability.  In one form, the LBV seems to traverse the H-R diagram at a constant luminosity, but with decreasing effective temperature/increasing radius and increase of 1-2 visual magnitudes during the outburst \\citep{puls08}.  This type of variability is known as S Doradus (SD) variability, and is seen as either a short SD phase, with a timescale less than 10 years, or a long SD phase, with a timescale greater than 20 years.  Another class, or perhaps a separate kind of eruption in the same star, is accompanied by both luminosity and mass increases, with ejection of several solar masses of material, and brightening by 1-3 bolometric magnitudes, and many visual magnitudes.  These stars, named the `supernova impostors' \\citep[e.g.,][]{smith11,vandyk12}, may have recurrences of these `giant eruptions' on timescales of centuries.  In our galaxy, the most well-known examples of the giant eruptions are the LBVs P Cyg and Eta Car.  Several candidates have also been detected in the Large Magellanic Cloud \\citep{wolf00} and external galaxies \\citep[e.g., M33,][]{clark12}.  A third kind of variation in the Luminous Blue Variables is the `microvariations' in luminosity, of a few tenths of a stellar magnitude.  These microvariations appear stochastic in nature, but have irregular periods of order of weeks to months \\citep{abolmasov11}.  See also \\cite{vink12} for a recent review of LBV properties and research questions.  The first and second types of variability recently have been proposed as an explanation for various features seen in supernova (SN) light curves.  For example, \\citet{kotak06} have proposed that SNe which show quasi-sinusoidal variation in their radio light curve have LBVs undergoing the SD type variations as progenitors.  As the mass-loss rates appear to change during a SD outburst \\citep[e.g.,][]{stahl01,stahl03}, the resulting circumstellar medium varies in density.  LBVs as progenitors of SNe have also been proposed as the progenitors of Type IIn SNe or superluminous SN \\citep[e.g.][]{smith10}.  The evolutionary status of LBVs is also uncertain.  It has long been accepted that massive O stars become LBVs before losing their outer layers to become Wolf-Rayet (WR) stars, and subsequently SN \\citep[e.g.,][]{chiosi86}.  Stars will lose mass through radiatively driven winds throughout these phases, and it has been shown that some LBV envelopes have enriched He abundances and evidence of CNO-cycle processing indicating that outer layers have been removed to expose material that had previously experienced hydrogen burning \\citep{najarro97,najarro01}.  However, with recent downward revisions of theoretical mass-loss rates, it seems that radiatively driven mass loss cannot remove enough material to match the observed masses of WR stars (about 20 \\msun).  If most or all massive stars end as WR stars, the most massive stars may need a series of giant eruption events during an LBV phase to reach the masses of typical WR stars \\citep{smith06}.  The LBVs are cooler than the core He-burning WR stars, which have been stripped of their outer layers, presumably during an earlier LBV phase.  The most luminous LBVs are hotter and brighter than the Humphreys-Davidson Limit \\citep[H-D Limit,][]{HD} in the H-R diagram (see Fig.\\,\\ref{evolvedmodels}), and may be on their first crossing of the H-R diagram in a shell H-burning or core He-burning phase.   At lower initial masses and luminosities, stars are found below the Humphreys-Davidson limit and across the width of the H-R diagram as blue, yellow, and red supergiants.  At present, the physical origin of the Humphreys-Davidson limit is not understood, and the highest initial mass which can evolve to the red supergiant region is determined observationally.  While much analytical work, atmosphere and wind modeling, and observational work has been assembled to understand and categorize LBV phenomena, and some explanations of the origin of LBV winds and outbursts have been proposed, to our knowledge very little research has been published in the refereed literature on the role of stellar pulsations in producing LBV microvariations or initiating outbursts or winds. In addition, a pulsational instability or some other hydrodynamic instability deep in the core must be responsible for initiating LBV \u00d4giant eruptions\u00d5.  Here we review the work of the group at Los Alamos National Laboratory from 1993 to the present that has before only been summarized briefly in abstracts or published in conference proceedings.\\footnote{Since the original publication of this paper, work in this area by Lovekin \\& Guzik has been submitted for publication in MNRAS.}  (These conference presentations are designated by asterisks in the References.)  We focus on linear stellar pulsation modeling using 1-D stellar evolution models including mass loss, and 1-D nonlinear hydrodynamic envelope models using the Dynstar Lagrangian stellar pulsation hydrodynamics code \\citep{ostlie82,ostlie90,cox83,cox90,coxostlie93} which includes a time-dependent convection treatment.  We also suggest some possible mechanisms for initiating LBV giant eruptions.  We hope that this review will motivate work by those who have more advanced computational tools to further investigate instabilities that may be responsible for LBV episodic mass loss and giant eruptions.  For a recent comprehensive overview of the properties of Eta Car and other known LBV stars, we recommend the review articles in the volume edited by Davidson and Humphreys, Eta Carinae and the Supernova Impostors, vol. 384 of the Springer Astrophysics and Space Science Library published in 2012. ",
        "conclusions": ""
    },
    "1402/1402.4521_arXiv.txt": {
        "abstract": "The global positioning system (GPS) is composed of thirty one satellites having atomic clocks with $10^{-15}$ accuracy on board and enables one to calibrate the primary standard for frequency on the ground. Using the fact that oscillators on the ground have been successfully stabilized with high accuracy by receiving radio waves emitted from the GPS satellites, we set a constraint on the strain amplitude of the gravitational wave background $h_{\\rm c}$.  We find that the GPS has already placed a meaningful constraint, and the constraint on the continuous component of gravitational waves is given as $h_{\\rm c}<4.8\\times 10^{-12}(1\\slash f)$ at $10^{-2}\\lesssim f \\lesssim 10^{0}$ Hz, for stabilized oscillators with ${\\Delta \\nu}\\slash{\\nu}\\simeq 10^{-12}$.  Thanks to the advantage of the Doppler tracking method, seismic oscillations do not affect the current constraint. Constraints on $h_c$ in the same frequency range from the velocity measurements by the lunar explorers in the Apollo mission are also derived. ",
        "introduction": "The existence of gravitational waves is predicted in general relativity and many modified gravity theories.  By observing gravitational waves (GWs), we can not only test the general relativity and modified gravity theories in strong gravitational fields, but also reveal the merger process of compact objects through the history of the universe and physics in the early universe such as inflation ({\\it e.g.} \\cite{2011PhRvL.107u1301G}).  In particular, GWs which originate from the cosmological inflation have almost scale-invariant spectrum and propagate freely since their generation, and thus the detection of such scale-invariant GWs can be considered as a direct proof of the existence of cosmological inflation. In order to prove the primordial GWs to be scale-invariant, we have to observe GWs with a wide frequency range and high sensitivity.  Precise measurements of primordial GWs at high frequencies will tell us about the thermal history of the early universe, which could not be reached in other ways. For example, it has been suggested that the amplitude of GWs at the frequency range $f \\gtrsim 10^{-2.5}$ Hz can be used to infer the reheating temperature \\cite{2008PhRvD..77l4001N}, and the epoch of quark-gluon phase transition and neutrino decoupling at lower frequency range $10^{-10} \\lesssim f \\lesssim 10^{-8}$ Hz \\cite{2009PhRvD..79j3501K}. If the universe underwent a strongly decelerating phase after inflation realized in quintessential inflation models \\cite{1999PhRvD..59f3505P}, inflationary GWs have a blue spectrum at higher frequency $f\\gtrsim 10^{-2}$ Hz \\cite{1999CQGra..16.2905G,2004CQGra..21.1761T}. Therefore, while the first detection of primordial GWs is expected through the CMB experiments which probe the waves at smallest frequencies ({\\it e.g.} \\cite{2000cils.book.....L}) it is significant to explore gravitational waves with a wide range of frequency.  In addition, measurements of GWs at high frequency will open a new window on the black hole merger events in the universe. Matsubayashi {\\it et al.} showed that the frequency of gravitational waves that originate from mergers of intermediate mass black holes (IMBHs) falls inside the range of $10^{-5}\\lesssim f \\lesssim 10^{1}$ Hz, and one can discover or set constraints on these merger events within several gigapersec from the Earth in future gravitational wave observation projects \\cite{2004ApJ...614..864M}.  In order to observe GWs, many kinds of gravitational wave detectors have been proposed and developed.  The Laser Interferometer Gravitational Wave Observatory (LIGO) has already set a strong constraint on the strain amplitude $h_{\\rm c}<4\\times 10^{-24}$ at the frequency of GWs around $10^{2}$ Hz \\cite{2009RPPh...72g6901A,2013PrPNP..68....1R}.  However, seismic oscillations interfere ground-based observations of gravitational waves at lower frequencies as $f\\lesssim 10$ Hz ({\\it see} \\cite{2001PhyU...44R...1G}).  Aiming at this frequency range, Torsion bar detectors such as TOBA have been developed. Ishidoshiro {\\it et al.} set a constraint on the strain amplitude of the continuous component of GWs as $h_{\\rm c}< 2\\times 10^{-9}$ at $f\\sim 0.2$Hz \\cite{2011PhRvL.106p1101I}.    At slightly lower frequency range, with some planetary explorers, constraints such as $h_{\\rm c}<1\\times 10^{-15}$ at $f=10^{-2}$ Hz \\cite{1995A&A...296...13B} and $h_{\\rm c}<2\\times 10^{-15}$ at $f=3\\times 10^{-4}$ Hz \\cite{2003ApJ...599..806A} have been set by the Doppler tracking method. However it is difficult to distinguish signals of GWs as the frequency of targeting GWs increases higher than $10^{-2}$ Hz due to the noise in the electric circuits of which are installed in the receiver in this method.  In this work, we consider the modulation of frequency of radio waves from GPS satellites by stationary/prompt GWs. GPS satellites emit precious, stable and high intensity radio waves in order to provide positional accuracy for the GPS in navigation. It is possible because the oscillators of GPS satellites are designed to be synchronized with atomic clocks which are loaded with the satellites. These oscillators are so stable that the fractional error of the frequency of their emitting radio waves is suppressed to $\\Delta \\nu \\slash \\nu<10^{-15}$ \\cite{2008.gps.standard}.  The GPS method is an application of the Doppler tracking method which has two advantages. First, GPS method can probe in the frequency range $f\\gtrsim 10^{-2}$ Hz where the conventional Doppler tracking methods can not reach because GPS satellites are much closer than the planetary explorers, while the distance is large enough for the noise due to seismic oscillations to be negligible as shown in the next section. Secondly, the radio waves from GPS satellites can be detected everywhere and everytime on the ground. This condition will be suitable to make cross correlation study and crucial to detect GWs from prompt events.  In this study, we set a constraint on the amplitude of GWs for $0.01 \\le f \\le 1$ Hz by using the Doppler tracking method with GPS disciplined oscillators (GPSDOs). The GPSDO is the stable oscillator with a quartz oscillator whose output is controlled to agree with the signals from GPS satellites \\cite{2008.gps.standard,2005ptti.conf..677L}.  To have GPSDOs operating with high stability, the amplitude of the continuous components of GWs should be small.  Recently, the frequency stability of GPSDOs for a short time interval about from seconds to a few hundred seconds has been reached to the level of ${\\Delta \\nu}\\slash{\\nu} \\simeq 10^{-12}$ \\cite{2008.gps.standard,2005ptti.conf..677L}. This short time stability of the frequency of GPSDOs enables us to set a constraint on the strain amplitude of the continuous GWs.  This paper is organized as follows. We estimate the effects of GWs on the measurements of the radio waves from GPS satellites and derive the constraint in \\S 2.  In our formulation, we adopt TT gauge to describe the GWs.  In \\S 3, we discuss implications of the result to the merger events of IMBHs and inflationary GWs. We conclude this paper in \\S 4. Throughout this paper, $\\nu$ and $f$ mean the frequency of the electro-magnetic waves and the GWs, respectively. The speed of light is denoted by $c$.  A dot represents a partial derivative respective to the physical time, {\\it i.e.} $\\dot{x}\\equiv {\\partial x}\\slash{\\partial t}$. ",
        "conclusions": "We set a constraint on the strain amplitude of the continuous component of GWs as $h_{\\rm c }<4.8\\times 10^{-12}\\left(1[{\\rm Hz}]\\slash f \\right)$ at the frequency range $10^{-2}\\lesssim f \\lesssim 1$ Hz with the radio waves emitted by GPS satellites in operation. Because the distance between GPS satellites and observers is order $10^{7}$ m, seismic oscillations do not affect the constraints on the strain amplitude.  The sensitivity to the GWs is limited to that of the A\\slash D converter on the GPS receiver at the frequency range."
    },
    "1402/1402.1903_arXiv.txt": {
        "abstract": "{% Disk accretion onto magnetic stars occurs in a variety of systems, including accreting neutron stars (with both high and low magnetic fields), white dwarfs, and protostars. We review some of the key physical processes in magnetosphere-disk interaction, highlighting the theoretical uncertainties. We also discuss some applications to the observations of accreting neutron star and protostellar systems, as well as possible connections to protoplanetary disks and exoplanets. } ",
        "introduction": "\\label{intro}  Disk accretion onto magnetic central objects occurs in a variety of astrophysical contexts, ranging from classical T Tauri stars (e.g., Bouvier et al.~2007) and cataclysmic variables (intermediate polars; e.g. Warner 2004), to accretion-powered X-ray pulsars (e.g. Lewin \\& van der Klis 2006). Such accretion has been studied extensively since the discovery of high-mass accreting X-ray pulsars in the 1970s (e.g., Pringle \\& Rees 1972). The basic picture of disk-magnetosphere interaction is well known. The stellar magnetic field disrupts the accretion flow at the magnetospheric boundary and funnels the plasma onto the polar caps of the star or ejects it to infinity (e.g., Pringle \\& Rees 1972).  The magnetosphere boundary is located where the magnetic and plasma stresses balance, \\be r_m =\\xi \\left({\\mu^4\\over GM\\dot M^2}\\right)^{1/7}, \\ee where $M$ and $\\mu$ are the mass and magnetic moment of the central object, $\\dot M$ is the mass accretion rate and $\\xi$ is a dimensionless constant of order 0.5-1. Roughly speaking, the funnel flow occurs when $r_m$ is less than the corotation radius $r_c$ (where the disk rotates at the same rate as the star). For $r_m\\go r_c$, centrifugal forces may lead to ejection of the accreting matter (``propeller'' effect; Illarionov \\& Sunyaev 1975).  However, beyond this simple ``Astro-101'' picture (with the estimate of $r_m$ given above), the situations are far more complicated.  Over the last several decades, numerous theoretical studies have been devoted to understanding the interaction between accretion disks and magnetized stars. Many different models have been developed.  Some examples are (see, e.g., Ghosh \\& Lamb 1979; Aly 1980; Lipunov \\& Shakura 1980; Anzer \\& B\\\"orner 1980,1983; Arons 1987,1993; Wang 1987,1995; Aly \\& Kuijpers 1990; Spruit \\& Taam 1990,1993; K\\\"onigl 1991; Shu et al.~1994,2000; van Ballegooijen 1994; Lovelace et al.~1995,1999; Li et al.~1996; Wickramasinge \\& R\\\"udiger 1996; Campbell 1997; Lai 1998,1999; Terquem \\& Papaloizou 2000; Shirakawa \\& Lai 2002a,b; Pfeiffer \\& Lai 2004; Uzdensky et al.~2002, Uzdensky 2004; Matt \\& Pudritz 2005; D'Angelo \\& Spruit 2010):  $\\bullet$ Ghosh \\& Lamb (1979) assumed that the stellar fields invade the disk over a large range of radii. Differential rotation between the star and the disk will cause the field lines to wind up. They argued that a steady state can be reached, where the winding-up of the field lines is balanced by dissipation, so that the field lines slip across the disk.  A problem with the GL picture is that in order to avoid too much winding, they need a very large (and unrealistic) magnetic diffusivity.  $\\bullet$ In fact, the accretion disk is a very good conductor. So it is possible that the stellar magnetic field cannot penetrate the disk except near the magnetosphere boundary. This kind of models were first explored by Arons, McKee \\& Pudritz (see Arons 1986).  The field lines are strongly pinched and a centrifugally driven wind may come out from the magnetosphere boundary.  These ideas were worked out and expanded in more details in the so-called x-wind model of Shu et al.~(1994), which has a rich set of phenomenology and applications.  $\\bullet$ Lovelace et al.~(1995) considered a model where the field lines connecting the star and disk becomes open because of the winding-up. So there is an inner magnetosphere corotating with the star but also an open region where wind may come out.  $\\bullet$ Matt \\& Pudritz (2005) developed a model where the possible role of stellar wind driven by accretion was emphasized..  \\begin{figure} \\centering \\includegraphics[width=12.5cm,clip]{f1.pdf} \\vskip -0.8truecm \\caption{Different theoretical models of magnetosphere-disk interaction (see text for discussion).} \\label{f1} \\end{figure}  In parallel to these theoretical studies, since 1990s there have been many numerical simulations, with increasing sophistication, on this problem (Hayashi et al.~1996,2000; Miller \\& Stone 1997; Goodson et al.~1997; Fendt \\& Elstner 2000; Matt et al.~2002; Romanova et al.~2003,2009). These simulations are playing an important role in elucidating the physics of magnetosphere-disk interaction in various astrophysical situations (see papers by Romanova and by Zanni in this volume).  In this paper, we will not discuss any particular models or simulations, which are covered by others. Instead, we will focus on several uncertain issues in the theory of magnetosphere-disk interaction, and will also discuss their possible applications and relevance in various astrophysical situations. ",
        "conclusions": ""
    },
    "1402/1402.1418_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  \\vskip 12pt  It would be bizarre if Supersymmetry \\cite{SUSY} were not to play a role in the Fundamental Interactions, since its local realization in Supergravity \\cite{SUGRA} and its completion in String Theory \\cite{strings} contain profound lessons on the links between gravity and the other forces. Yet, the apparent lack of super--partners of the known particles has already survived the first round of experiments at LHC \\cite{LHC}. Under these circumstances, the standard recipe to keep Supersymmetry alive is to try and make all putative partners heavy enough, appealing to some mechanism of supersymmetry breaking \\cite{SUSY_breaking}. However, arriving at a fully satisfactory scenario of this type remains a key challenge in current attempts to combine the Standard Model of Particle Physics with gravity, despite decades of intense effort and a number of important results.  A proper understanding of supersymmetry breaking cannot forego a detailed grasp of general matter couplings in Supergravity, which has long been available for models with ${\\cal N}=1$ supersymmetry in four dimensions \\cite{cfgv} but is not nearly as complete in the higher--dimensional settings that are so important in String Theory. As a result, perhaps, non--supersymmetric string compactifications are generally expected to spell trouble for the vacuum state and have been explored only to a limited extent. Still, there is a higher--dimensional setting that was identified long ago and stands out for its relative simplicity and rigidity. This is ``brane supersymmetry breaking'' (BSB) \\cite{bsb}, which presents itself in some orientifold vacua \\cite{orientifolds} where Ramond--Ramond (RR) charge neutrality requires the simultaneous presence, in the vacuum, of combinations of branes and orientifolds that are not mutually BPS and yet do not introduce tree--level tachyon instabilities. These extended objects leave behind a distinctive mark, a steep exponential potential proportional to their overall tension, whose lack of local minima excludes from the outset maximally symmetric geometries, and flat space in particular, for these systems. The phenomenon finds its simplest manifestation in the ten--dimensional Sugimoto model in \\cite{bsb}, where supersymmetry appears exact, at tree level, insofar as the closed spectrum is concerned, but is actually non-linearly realized due to the non-supersymmetric brane configuration, whose modes include a singlet spinor that plays the role of a goldstino \\cite{10d_bsb_couplings}. The nine dimensional vacuum geometry realizing the largest symmetry allowed for this system was presented by Dudas and Mourad in \\cite{dm_vacuum}.  The ten--dimensional Sugimoto model in \\cite{bsb} admits an intriguing spatially flat cosmological solution where the dilaton exhibits a striking behavior. The solution has a long history \\cite{lm,exp_sol,dm_vacuum}, but its main lesson was appreciated only recently \\cite{dks}. A minimally coupled scalar field ought to possess the two distinct options of emerging from an initial singularity while \\emph{descending} or \\emph{climbing} mild exponential potentials of the type \\be {V}(\\varphi) \\ = \\ {V}_0 \\, e^{\\,2\\,\\gamma\\, \\varphi} \\ , \\label{potonexp_intro} \\ee and this is indeed the case. However, \\emph{for $\\gamma$ large enough only the climbing behavior becomes possible}, and this what we shall refer to as the climbing phenomenon. In particular, with the convenient non--canonical normalization for $\\varphi$ described in the following sections the climbing behavior sets in at $\\gamma=1$ for all space--time dimensions $d$.  A striking feature of ``brane supersymmetry breaking'' is that its potential lies precisely at the critical point $\\gamma=1$ \\cite{dks}. Moreover, this property continues to hold for $d<10$ for a special combination of the dilaton and the breathing mode, and under the assumption that this field dominates the early dynamics other branes present in String Theory \\cite{branescan,branesugimoto} can contribute milder exponential terms \\cite{dks,fss,as13} that could have injected inflation \\cite{inflation} after the initial climbing phase. One is thus led to consider the class of potentials \\cite{dks,dkps} \\be {V}(\\varphi) \\ = \\ {V}_0 \\left( e^{\\,2\\,\\varphi} \\ + \\ e^{\\,2\\,\\gamma\\, \\varphi} \\right) \\  \\label{potwoexp_intro} \\ee in four dimensions, where $\\gamma< 1/\\sqrt{3}$ in order to allow for the onset of inflation following an initial climbing phase, with the eventual goal of comparing their implications with the low--$\\ell$ end of the CMB angular power spectrum.  An enticing feature of the climbing phenomenon is that it links two apparently unrelated problems, the breaking of Supersymmetry and the onset of inflation. We are well aware of the limitations of the simple potentials of eq.~\\eqref{potwoexp_intro}, which result for one matter in tensor--to--scalar ratios that are too large \\cite{inflation}, but we can anticipate that the key phenomenon that we shall come to can be traced to the behavior of the scalar field near the ``hard'' exponential, the main datum that String Theory and ``brane supersymmetry breaking'' contribute to this discussion, so that it occurs in a variety of more realistic potentials, some of which will be touched upon in the following.  Some features of the scalar dynamics in the potential \\eqref{potwoexp_intro} can be readily anticipated. To begin with, the ``hard'' potential forces the field to emerge from an initial singularity while climbing up from large negative values of $\\varphi$, and this early phase is essentially driven by the mild exponential. Notice that the climbing phenomenon constrains the choice of initial conditions, a nice feature for a theory of inflation, so that a single parameter is left in this case, a constant $\\varphi_0$ that determines to which extent the scalar feels the ``hard'' exponential. Depending on the value of $\\varphi_0$, the reversal that opens the descent can be more or less abrupt, but lo and behold one would expect it to bring along a spurt of exponential expansion for the Universe, before the actual inflationary phase sets in. This generic feature was readily recognized and actually served as a motivation for the analysis in \\cite{dks}, but turning into a quantitative prediction for the spectrum of scalar perturbations proved rather difficult, so much so that it failed to emerge in \\cite{dkps}. As we shall see, a pre--inflationary peak does show up in the power spectrum of scalar perturbations that, as discussed in \\cite{dkps}, experiences in general a wide infrared depression before merging with an almost scale invariant profile ${\\cal P}_\\zeta(k) \\sim k^{n_s-1}$ as the dynamical evolution finally approaches the Lucchin--Matarrese (LM) attractor \\cite{lm}. The reversal becomes less abrupt for lower values of $\\varphi_0$, while the peak grows in size, until it eventually turns into the typical feature well described in \\cite{destri}. This signals the approach to slow roll in more conventional models, and actually shows up, in the presence of a mild exponential alone, in the whole range of values for $\\varphi_0$ that we have explored. As expected, however, the reversal of the scalar motion leaves no signs in the corresponding spectrum of tensor perturbations, consistently with the analysis in \\cite{dkps}.  It is natural to inquire whether this type of behavior could have some bearing on our current understanding of the CMB angular power spectrum. We shall qualify under which assumptions this might well be the case, and we shall also manage to vindicate, to some extent at least, our expectation via a detailed, if preliminary, comparison with CMB data. As we have anticipated, the region of interest is the low--$\\ell$ tail of the angular power spectrum, where some anomalies with respect to the $\\Lambda$CDM setting have long been spotted but where cosmic variance, which reflects our very special observation site for the Universe, adds more than one word of caution to any attempt to interpret their actual meaning.  The main low--$\\ell$ anomaly in the CMB angular power spectrum is quadrupole reduction, and it is large enough to go unnoticed. Indeed, a number of works recently touched upon the subject from different viewpoints \\cite{dkps,quadrupole_red}, and they include a detailed re--analysis of the cosmic mask \\cite{gruppuso}. Interestingly one can argue, on rather general grounds, that \\emph{quadrupole reduction accompanies naturally the emergence from an initial singularity}. Moreover, for low multipoles $\\ell \\lesssim 35$ the actual CMB observable, the angular power spectrum, is determined essentially by a Bessel--like transform \\cite{mukhanov_slow} that follows closely the power spectrum, so that in principle features of the former can reflect themselves in similar features of the latter, and vice versa.  In practice, however, this type of correspondence requires an additional assumption that, if correct, would make, by itself, the whole story quite interesting. The assumption, which is not implausible numerically, posits that the largest wavelengths entering the horizon at the present epoch are essentially those that exited around the onset of the inflationary phase. Or, if you will, that the low--$\\ell$ CMB data are opening in front of our instruments a small window on the onset of inflation, the very phenomenon that is usually advocated to explain the apparent flatness and homogeneity of our Universe and also explains naturally the slight tilt of the CMB power spectrum \\cite{cm} that was recently confirmed to high precision by PLANCK \\cite{PLANCK}. Working within this assumption, we shall begin to explore how far one can go in relating the available WMAP9 raw data \\cite{wmap9} to the models at stake. We shall explore, to this end, the first 31 multipoles starting from the quadrupole, for a range of values of $\\varphi_0$ that encompasses the emergence of the pre--inflationary peak, its growth and its eventual coalescence into the attractor profile.  As we shall see, the data are apparently not insensitive to the pre--inflationary peak, since centering it around $\\ell = 5$ brings about a noticeable reduction of the $\\chi^2$ value by two or three units. Amusingly the agreement improves, as we shall see, for values of $\\gamma$ that lie below $0.08$, the choice that would tune the large -- $k$ behavior of ${\\cal P}_\\zeta(k)$ with the observed spectral index $n_s \\approx 0.96$. This result resonates with a key indication of the PLANCK experiment, which favors generically concave inflationary potentials \\cite{concave}. As we have anticipated, the pre--inflationary peak draws its origin from the region where the ``hard'' exponential begins to dominate and only the nearby behavior, which is naturally flatter in a concave potential, should play a role. We shall also vindicate this claim by displaying some power spectra computed directly in Starobinsky--like potentials \\cite{starobinsky} \\be {V}_S(\\varphi) \\ = \\ {V}_0\\,\\left[ \\left(1- e^{\\,-\\,\\gamma\\,(\\varphi+\\Delta)}\\right)^2 \\ + \\ e^{\\,2\\,\\varphi}\\right] \\label{starobinsky_intro} \\ee that terminate on the same hard exponential, which possess very similar qualitative features. The Starobinsky potentials have aroused some interest lately in connection with Supergravity \\cite{starobinsky_supergravity}, and are not foreign, in principle, to ``brane supersymmetry breaking'', if quantum corrections are taken into account. Interestingly, as we shall see, the comparison with CMB raw data favors scenarios of this type that appear to fit well within perturbative String Theory. The string coupling is in fact sized by $e^\\varphi$, a quantity that stays well below one for the choices of $\\varphi_0$ that result in better fits. We plan to return soon to a more detailed comparison with the CMB, modifying the standard $\\Lambda$CDM setup to allow for quadrupole reduction and a pre--inflationary peak, the key features suggested by this class of models \\cite{gnks}.  To reiterate, among a multitude of available vacua, String Theory suggests some peculiar options related to orientifold models where supersymmetry is broken at the string scale \\cite{bsb}. Orientifold models generally allow a wide range of values for the string scale \\cite{aadd}, all compatible with the standard values of Newton's constant and of gauge couplings, which includes the scales typically associated with inflation. And, as we have stressed, these two scales are linked in the simplest realizations of ``brane supersymmetry breaking'' with a ``climbing scalar''. A common origin for the two phenomena of supersymmetry breaking and inflation would represent an economy of principles, and our results can perhaps serve as a motivation in this respect, although they do not arise generically but only within a specific class of string vacua.  The structure of the paper is as follows. In Section \\ref{sec:climbing} we review the climbing phenomenon starting from the one--exponential case, stressing its generality and illustrating its realization in the relatively simple class of potentials of eq.~\\eqref{potwoexp_intro} and in the richer class of potentials of eq.~\\eqref{starobinsky_intro}. In Section \\ref{sec:powerspectrum}, which focusses on power spectra of scalar perturbations, we pinpoint the origin of the pre--inflationary climbing peak and we illustrate its dependence on $\\varphi_0$. We also discuss briefly, for the sake of comparison, corresponding spectra of tensor perturbations for the same range of values for $\\varphi_0$. In Section \\ref{sec:observables} we move some first steps toward a quantitative comparison with the CMB, insofar as the first 30 multipoles or so are concerned. Finally, in Section \\ref{sec:conclusions} we briefly summarize our main conclusions and some perspectives for future research along these lines, while the Appendix elaborates on the links between the two--exponential potentials \\eqref{potwoexp_intro} and String Theory. \\vskip 24pt ",
        "conclusions": "\\label{sec:conclusions} \\vskip 12pt  This work builds on two main inputs. The first, drawn from String Theory \\cite{strings}, is the existence of a class of orientifold vacua \\cite{orientifolds} with ``brane supersymmetry breaking'' \\cite{bsb}. In these models, which admit no maximally symmetric vacuum geometries, supersymmetry is broken at the string scale and is non-linearly realized in the low--energy supergravity \\cite{10d_bsb_couplings}, which includes a ``hard'' exponential potential, but no tachyon excitations are present at tree level. The second input is drawn from a striking feature of the spatially flat cosmologies allowed by the corresponding low--energy supergravities. These involve a scalar field that emerges from an initial singularity with no other option than climbing up the steep exponential potential \\cite{exp_sol,dks}. The process comes to an end at a turning point, while other branes of String Theory \\cite{branescan,branesugimoto} can give rise, in principle, to milder exponentials \\cite{fss,as13} that can force an inflationary phase during the ensuing descent. Therefore, in this setting the breaking of supersymmetry at high scales can provide a rationale for inflation to begin.  Intuitively, as the climbing phase ends, in reverting its motion the scalar ought to bring along a spurt of exponential expansion for the Universe and a corresponding peak in the power spectrum of scalar perturbations. Here we showed that this is indeed the case for scalar perturbations, while tensor perturbations, which depend essentially on the scale factor alone, do not exhibit a similar phenomenon. It took some effort to quantify this expectation, but we have provided ample evidence for it here, and we have also started to compare these findings with the low--$\\ell$ tail of the CMB angular power spectrum. \\emph{Leaving aside cosmic variance, a fair summary of our findings is that the low--$\\ell$ WMAP9 raw data tend to favor slightly scenarios of this type with respect to the attractor power spectrum underlying the standard $\\Lambda$CDM setup}. The most interesting aspect of the whole setting, however, is perhaps the main assumption on which the comparison rests. This posits an essentially direct correspondence between the largest wavelengths entering the cosmic horizon at the present epoch and those that exited at the onset of inflation. If true, it would translate into the enticing perspective of drawing from the low--$\\ell$ tail of the CMB power spectrum some information on the very early Universe. \\begin{figure}[h] \\begin{center}$ \\begin{array}{cccc} \\epsfig{file=starobinsky_1.eps, height=1.2in, width=1.2in}& \\qquad \\epsfig{file=starobinsky_2.eps, height=1.2in, width=1.2in}& \\qquad \\epsfig{file=starobinsky_3.eps, height=1.2in, width=1.2in}& \\qquad \\epsfig{file=starobinsky_4.eps, height=1.2in, width=1.2in} \\end{array}$ \\end{center} \\caption{\\small Power spectra of scalar perturbations computed, in cosmic time, for the Starobinsky--like potentials of eq.~\\eqref{starobinsky}, shown here with an arbitrary normalization but with the parameters adjusted in order to guarantee about 60 $e$--folds of inflation and a fair portion of them with $n_s \\simeq 0.96$. The slopes of the dotted lines reflect the range of values for $n_s$ that are consistent with current observations.} \\label{fig:starobinsky_power} \\end{figure}  All we have shown so far rests on the two--exponential potentials of eq.~\\eqref{potwoexp}, which are relatively simple to analyze but are clearly incomplete in several respects. As we have stressed, however, the pre--inflationary peaks reflect the local behavior in the region where the scalar reverts its motion, and we have also seen that the agreement with WMAP9 data improves slightly as the parameter $\\gamma$ of the two--exponential models is reduced below 0.08, consistently with the expected preference for concave potentials. We can actually conclude with a brief discussion of some results obtained directly for the Starobinsky--like potentials of eq.~\\eqref{starobinsky}, with parameters adjusted so as to grant about 60 $e$--folds of inflation and a fair portion of them with $n_s \\simeq 0.96$. The computation was more subtle since, for one matter, these potentials are negligibly small near the minimum so that we could not resort to the convenient gauge \\eqref{gauge}, and moreover the initial values were to be set very close to the singularity. Still, the ``lsode'' method available in Maple worked rather efficiently, and the power spectra displayed in fig.~\\ref{fig:starobinsky_power} clearly vindicate our claims. As in the two--exponential models, tuning the impact with the hard exponential a slow growth leaves way to the birth of a pre--inflationary peak and to its eventual merge with an approximately scale--invariant profile. The curves collected in the figure are indeed very similar to those displayed in Section \\ref{sec:powerspectrum}, but turn more gradually toward the slope dictated by the observed spectral index $n_s \\simeq 0.96$, which we used as an input. We did not perform a systematic analysis as for the two--exponential case, but for example the power spectra displayed in fig.~\\ref{fig:starobinsky_power} yield for $\\chi^2_{\\rm min}/DOF$ the four values 0.738, 0.724, 0.783 and 0.816. These results are along the lines of those displayed in fig.~\\ref{fig:angular_fit_double}, and in particular the second reaches the minimum of the curve for $\\gamma=0.04$. Let us stress, to conclude, that only a small portion of the power spectra centered around the pre--inflationary peak played a role in our low--$\\ell$ analysis via eq.~\\eqref{bessel}. There is thus room, in principle, for power spectra of this type to be compatible with the $\\Lambda$CDM analysis of higher multipoles, while additional small features in the potential could well account for the other oscillations that are apparently present for $\\ell \\lesssim 35$ in fig.~\\ref{fig:attractor-optimalCl}.  Two very legitimate objections could be raised, within String Theory, against our analysis. The first concerns the values of the string coupling that accompany these phenomena, and in this case we could provide an encouraging answer, since the scenarios that are apparently preferred in the comparison with the CMB rest on relatively small values of the string coupling. There is a second objection, however, on which we have nothing definite to say, now as in the past. It has to do with curvature corrections, which are ubiquitous in String Theory and are expected to dominate near an initial singularity, casting doubts on a low--energy analysis of the climbing phenomenon. They were examined, insofar as possible, at low orders in \\cite{cd}, but it is fair to state that at present we do not understand if and how they could be under control. Time and more detailed studies will tell whether these considerations can find a more rigorous origin in String Theory, and whether the PLANCK data to be soon released and a more refined data analysis taking into account wider portions of the CMB power spectrum \\cite{gnks} will confirm some encouraging clues that have emerged from the present work.  \\vskip 24pt"
    },
    "1402/1402.1160.txt": {
        "abstract": "{The origin of protostellar jets remains a major open question in star formation. Magneto-hydrodynamical (MHD) disk winds are an important mechanism to consider, because they would have a significant impact on planet formation and migration.} {We wish to test the origins proposed for the extended hot H$_{2}$ at 2000~K around the atomic jet from the  T Tauri star DG~Tau, in order to constrain the wide-angle wind structure and the possible presence of an MHD disk wind in this prototypical source.} {We present spectro-imaging observations of the DG~Tau jet in H$_{2}$ 1-0 S(1)  with 0\\farcs12 angular resolution, obtained with SINFONI/VLT. Thanks to spatial deconvolution by the PSF and to careful correction for wavelength calibration and for uneven slit illumination (to within a few \\kms), we performed a thorough analysis and modeled the morphology and kinematics. We also compared our results with studies in [Fe II], [O I], and  FUV-pumped H$_2$. Absolute flux calibration yields the H$_2$ column/volume density and emission surface, and narrows down possible shock conditions.} {The limb-brightened H$_{2}$ 1-0 S(1) emission in the blue lobe is strikingly similar to FUV-pumped H$_2$ imaged 6 yr later, confirming that they trace the same hot gas and setting an upper limit $<12$\\kms\\ on any expansion proper motion. The wide-angle rims are at lower blueshifts (between -5 and 0 \\kms) than probed by narrow long-slit spectra. We confirm that they extend to larger angle and to lower speed the onion-like velocity structure observed in optical atomic lines. The latter is shown to be steady over $\\ge 4$ yr but undetected in [Fe II] by SINFONI, probably due to strong iron depletion.  The rim thickness $\\le 14$ AU rules out excitation by C-type shocks, and J-type shock speeds are constrained to $\\simeq$ 10 \\kms.} {We find that explaining the H$_2$ 1-0 S(1) wide-angle emission with a shocked layer requires either a recent outburst (15 yr) into a pre-existing ambient outflow or an excessive wind mass flux. A slow photoevaporative wind from the dense irradiated disk surface and an MHD disk wind heated by ambipolar diffusion seem to be more promising and need to be modeled in more detail. Better observational constraints on proper motion and rim thickness would also be crucial for clarifying the origin of this structure.} %{Context}{Aims}{Methods}{Results}{Conclusions} % context heading (optional) % aims heading (mandatory) % methods heading (mandatory) % results heading (mandatory) % conclusions heading (optional), leave it empty if necessary % 5 {} token are mandatory ",
        "introduction": "Jets from young stars are ubiquitous in star formation, and their link with accretion processes is well established \\citep{Cabrit1990,Hartigan1995}. They are launched from the inner regions of the star-disk system, the same disk that may ultimately give rise to a planetary system. Although a number of models exist for their generation (see, for example, the reviews of \\citet{Ferreira2006}, \\citet{Shang2007} and \\citet{Pudritz2007}), the exact mechanism by which mass is ejected from accreting systems and collimated into jets is an open problem. It is thought that the same basic magnetohydrodynamic (MHD) mechanism is responsible for  launching jets in objects as diverse as brown dwarfs \\citep{Whelan2005}, post-AGB stars (e.g., \\citet{Garcia-Diaz2008}, \\citet{Garcia-Segura2005}), compact objects in binary systems, and perhaps even active galactic nuclei. %It is also important for telluric planet formation models as they have implications on inner disk physics. Because of their proximity and rich set of diagnostic emission lines, jets from young stellar objects represent a particularly useful test bed for the MHD jet-launching paradigm \\citep[see the reviews of][]{Bally2007,Ray2007}. %(see the reviews of \\citep[][]{Bally2007} and \\cite{Ray2007}).  %Studies of H2 in young stars Molecular hydrogen is the primary gas constituent in the circumstellar environment of young stars.  The $v$=1-0 S(1) transition of  H$_2$ at 2.12 $\\mu$m  is often detected in embedded protostars, with the emission formed primarily in shock-excited collimated outflows \\citep[e.g.,][]{Davis2011,Davis2002}. The presence of H$_2$ in jets might be an indication that MHD ejection operates out to disk radii $\\ge 0.2-1$AU where molecules can survive against dissociation \\citep{Safier1993,Panoglou2012}. If such extended MHD disk winds persist until the later planet building phase, they could maintain fast accretion across the disk ``dead-zone\" where ionization in the midplane is too low to activate MHD turbulence \\citep{Bai2013b}. The associated magnetic fields would have strong consequences for the  formation and migration of planets \\citep{Fromang2005,Bai2013}. The H$_2$ emission from warm (T=1000-3000 K) gas was recently detected within 100~AU of more evolved T Tauri stars (TTSs) and Herbig Ae/Be stars, both in the near-infrared and FUV domains \\citep[e.g.,][]{Valenti2000, Ardila2002, Bary2003, Carmona2011}. When studied at high spectral resolution, these lines split into two populations: most show narrow profiles centered on the stellar velocity, suggesting an origin in a warm disk atmosphere (quiescent H$_2$), but $\\simeq$ 30 \\% of them display profiles blueshifted by 10-30 \\kms, sometimes with blue wings extending up to -100 \\kms, suggesting formation in an outflow \\citep{Takami2004, Herczeg2006, Greene2010, Carmona2011, France2012}. In particular, long-slit spectra of DG Tau provide evidence of a warm H$_2$ outflow at $\\simeq$ -20 \\kms, shifted by $\\simeq$ 30 AU along the blue jet direction and with an estimated width $\\simeq$ 80 AU \\citep{Takami2004, Schneider2013a}. %Independent indications for small-scale low-velocity molecular outflows in the vicinity of TTSs have been also reported in CO lines by \\citet{Pontopiddan2011} in the near-IR, and by \\citet{Pety2006,Launhardt2009} in the millimetric range. %Determining the origin of H$_2$ emission (disk, outflow or residual envelope) is critical for our understanding of the gaseous component in protoplanetary disks and its evolution. \\citet{Beck2008} conducted a 2D spectro-imaging study of H$_2$ rovibrational lines in six CTTs known to be associated with large scale atomic jets, including DG Tau, and found spatially extended emission in all cases, with sizes up to 200 AU. While the data are not flux calibrated, relative level populations among K-band H$_2$ lines indicate excitation temperatures in the range 1800 -- 2300 K. From various arguments, \\citet{Takami2004} and \\citet{Beck2008} argue that the H$_2$ properties are in general not well explained by FUV or Xray irradiation and are most consistent with shock-excited emission from the inner regions of the atomic jets or from wider-angle winds encompassing these flows. \\citet{Takami2004} also mention ambipolar diffusion on the outer molecular streamlines of an MHD disk wind as an alternative heating mechanism for the warm H$_2$ outflow in DG Tau. %In an independent  FUV imaging study published while this work was being revised, \\citet{Schneider2013b} also favor a molecular MHD disk wind, but heated by shocks rather than by ambipolar diffusion. In HL Tau, on the other hand, \\citet{Takami2007} interpret the V-shaped H$_2$ emission encompassing the atomic jet as a shocked cavity driven into the envelope by an unseen wide-angle wind. In any case,  extended H$_2$ emission in TTS clearly holds key information on their wide-angle wind structure, the radial extension of any MHD disk wind present, and its impact on protoplanetary disk physics. %The detection of small-scale H$_2$ outflows in the vicinity of accreting T Tauri stars raises a few challenging questions. First of all, how is the H$_2$ excited? shocks (either in the outflow or against a static envelope/disk), UV fluorescence pumped by FUV or Ly$\\alpha$ photons, UV and X-ray heating from the star or the accretion shock, and ambipolar diffusion heating in a magnetic wind. Secondly, which is the physical mechanism responsible for the ejection: a thermal evaporative process similar to the one invoked for the low velocity atomic outflow signatures \\citep{Ercolano} or the same magneto-centrifugal process thought to be responsible for the larger scales atomic jets ? Indeed, \\cite{Panoglou2012} recently computed the detailed physico-chemical structure for steady MHD disk wind solutions and showed that the gas retains a significant H$_2$ abundance for streamlines launched beyond r$_0$=1 AU.  %To date, one of the most convincing case of a small-scale H$_2$ outflow is the one detected in the close vicinity of DG Tauri.  This object is a very active T Tauri star, with an accretion rate of 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$ \\citep{}, associated with a bright small scale atomic jet \\citep{Dougados2000, Bacciotti2000}. Because the central source suffers only moderate extinction, unlike younger Class I sources, this jet has been the target of many studies aimed at constrainning the launching mechanism of jets (e.g. \\citet{Agra2011} and references therein).  In an effort to extract this information, we present in this paper an in-depth analysis of  H$_2$ ro-vibrational 1-0 S(1) emission associated with DG~Tauri, using flux-calibrated spectro-imaging K-band observations with 0\\farcs12 resolution conducted with SINFONI/VLT. This work complements our detailed study of the atomic flow component and accretion rate in DG Tau based on [Fe II] lines in H-band and Br${\\gamma}$ in K-band, observed at the same epoch with the same instrument \\citep[][hereafter, Paper~I]{Agra-Amboage2011}.  Our analysis improves on that in \\citet{Beck2008} in several ways: We apply spatial deconvolution to obtain a sharper view of the brightness distribution, and perform careful correction for wavelength calibration and uneven-slit illumination to retrieve radial velocities along and across the jet axis down to a precision of a few \\kms. This allows us to conduct a detailed comparison with the spatio-kinematic structure of the atomic flow component and FUV H$_2$ emission in DG Tau, and to constrain both the proper motion and the 3D velocity field of warm H$_2$, including rotation. In addition, the absolute flux-calibration of the 1-0 S(1) line is used to infer column densities and emitting surfaces and, together with its ratio to  2-1 S(1)  at 2.25 $\\mu$m, to restrict the possible shock parameters for H$_2$ excitation.  We then combine this new information to critically re-examine the main scenarii that have been invoked for the origin of wide-angle  H$_2$ emission in DG Tau and other TTS, namely: (1) an irradiated disk/envelope, (2) a molecular MHD disk wind heated by ambipolar diffusion, (3) a forward shock sweeping up the envelope, (4) a reverse shock driven back into a wide-angle molecular wind. %The analysis focusses on the brightest K-band transition 1-0 S(1) at 2.12 $\\mu$m, but also uses its ratio to the 2-1 S(1) transition at 2.25 $\\mu$m to investigate possible excitation gradients and allowed shock parameters}.  %{\\bf While this paper was being revised, a 0\\farcs1 resolution image of FUV H$_2$ emission in DG Tau was obtained by \\citet{Schneider2013b} with HST using differential broad-band imaging. They independently reach the same conclusions as we do on a number (but not all) points. Comparison with our data will allow us to constrain the expansion proper motion of the H$_2$ cavity.} The paper is organized as follows: The observations, data reduction procedure, and centroid velocity measurements are described  in \\S 2.  In \\S 3 , we present the new results brought by our observations on the morphology, kinematics, and column density of the wide-angle H$_2$, and compare with the atomic  and FUV H$_2$ emission properties. In \\S 4, we clarify the shock parameters that could explain the H$_2$ surface brightness and line ratio, we explore the possible 3D velocity fields, and we discuss the strengths and weaknesses of various proposed origins for the wide-angle H$_2$. Our conclusions are summarized in \\S 5. Throughout the paper we assume a distance of 140 pc to the object.  %{\\bf We also explore the constraints on the true geometry and velocity field of the wide-angle H$_2$ set by deconvolved emission maps and radial velocities. We then combine these constraints to} critically re-examine {\\bf in the case of DG Tau} the main scenarii that have been invoked for the origin of {\\bf extended}  H$_2$ emission in TTS, namely: (1) an irradiated disc/envelope, (2) a molecular MHD disk wind heated {\\bf by ambipolar diffusion}, (3) a shock {\\bf caused by the interaction of a wide-angle wind with the disc/envelope. We show that all of these options meet some challenges in reproducing all of the observed features}. ",
        "conclusions": "\\label{sec:conclusions}  We have conducted a thorough analysis of H$_{2}$ 1-0 S(1) spectro-imaging data towards the classical T Tauri star DG Tau, with an angular resolution of 0\\farcs12 and a precision in velocity centroids reaching down to a few \\kms. We compared the morphology and kinematics with simultaneous data in [Fe II] \\citep{Agra-Amboage2011}, near-contemporaneous spectra in [O I] \\citep[][]{Coffey2007}, and spectra and images in FUV-pumped H$_2$ obtained 6 years later \\citep{Schneider2013a, Schneider2013b}. The absolute line brightness and the 2-1 S(1) / 1-0 S(1) ratio were used to constrain shock conditions, based on state-of-the-art shock models \\citep{Kristensen2008}. Projection effects in radial velocities were modeled to constrain the intrinsic velocity field of the emitting gas. This new information was used to re-examine proposed origins for the wide-angle H$_2$ emission in the blue lobe of DG Tau and similar sources. Our main conclusions, and our suggestions for further work, are the following: \\begin{itemize}  \\item The faint arc of H$_2$ 1-0 S(1) emission in the red lobe traces slow material at +5 \\kms\\ in the wings of a large bowshock prominent in [Fe~II]. In contrast, the limb-brightened H$_2$ 1-0 S(1) at the base of the blue lobe is much wider than the [Fe II] jet; it appears filled-in by the wide-angle atomic wind at low velocities seen in [O I] and [S II], that is, steady over $\\ge 4$yr but invisible in [Fe II] probably because of strong iron depletion. Most of the wide-angle H$_{2}$ 1-0 S(1)  is at low blueshifts (between -5 and 0 \\kms). Hence, we confirm previous claims that the H$_2$ blue lobe in DG Tau appears as a  slower, wider extension of the nested velocity structure previously observed in atomic lines \\citep{Takami2004, Schneider2013b}. The H$_2$ may be rotating about the jet axis in the same sense as the disk and atomic flow, with $V_\\phi \\le 7$ \\kms.  \\item The blue lobe morphology in H$_{2}$ 1-0 S(1) is strikingly similar to the FUV image \\citep{Schneider2013b}, confirming that they trace the same hot gas at 2000~K and setting upper limits on the rim thickness $\\le 14$~AU and on any expansion proper motion $V_{\\rm exp} < 12$\\kms. The face-on surface brightness in 1-0 S(1) is $\\simeq 3 \\times10^{-3}$ erg s$^{-1}$cm$^{-2}$ sr$^{-1}$, and the 2-1 S(1) / 1-0 S(1) ratio is uniform at our resolution. If due to shock-heating, they require a J-shock at 10 \\kms\\ into $10^6$ \\cm\\ gas (C-shocks would predict excessive rim thickness). Modeling of the radial velocities limits the gas motion perpendicular to the walls  to $-5 \\le V_\\perp \\le +5$ \\kms, unless there is a comparable parallel component in which case $V_\\perp$ can reach  20\\kms. %Besides, we discussed the three main scenarios proposed in order to explain the origin of this molecular emission. We summarize the main conclusions in the following: %\\begin{itemize}  \\item A photoevaporating disk atmosphere irradiated by the strong stellar FUV excess and Ly$\\alpha$ from jet shocks seems able to explain the excitation and brightness on observed scales, as well as the low blueshifts over most of the area. One caveat is the high density in the upper disk atmosphere of \\nH = $10^{7}$ \\cm\\  suggested by steady PDR models. Time-dependent PDR models including disc photoevaporation may alleviate this problem and would be highly desirable.  \\item A molecular MHD disc wind launched around 10 AU and heated by ambipolar diffusion is another promising scenario that can reproduce the observed morphology, kinematics (including limits on rotation), mass flux rate, and temperature  \\citep[see][]{Panoglou2012}. Synthetic predictions in H$_2$ 1-0 S(1) taking into account non-LTE effects, the full 3D velocity field, and a range of streamlines are necessary to fully test this hypothesis and will be the subject of future work. %The main problem in this case, is the lack of rotation evidences in our H$_{2}$ emission.% Predicted PV diagrams from detailed MHD disc wind models are necessary to fully test this scenario.  \\item A cavity of shocked ambient swept-up gas does not reproduce the H$_2$ 1-0 S(1) excitation and blueshifted centroids unless it is very young (15 yr) and expands into a pre-existing dense molecular wind of unclear origin moving at $\\simeq$15 \\kms\\ along the cavity walls. A crucial check of this scenario would be to detect the predicted proper motion $V_{\\rm exp} \\simeq 10$\\kms, and the signs of deceleration or recurrence expected for such an episodic phenomenon.  \\item A reverse shock into a wide-angle molecular wind could readily explain the excitation, low blueshifts, and apparent stationarity of the H$_2$ cavity, but the required one-sided wind mass flux of $\\ge 6.7 \\times 10^{-8} M_\\odot$ yr$^{-1}$ appears excessive compared to current accretion rates in DG Tau. \\end{itemize} We point out that the (J-type) shock scenarios predicts a FWHM of only 0.1 AU for the H$_2$ 1-0 S(1) emitting layer. This is to be compared with a predicted thickness of 7 AU for a static PDR ($N_H \\simeq 10^{21}$ cm$^{-2}$ and \\nH $\\simeq 10^7$ \\cm) and 4--14 AU for an MHD disk wind. Spatially resolving the rims of the H$_2$ cavity would thus bring useful insight into its origin. We also caution that we have considered simplistic scenarios where only one heating mechanism is operating at a time. A combination of shock and FUV irradiation might alleviate the caveats encountered when each is considered separately.  %In summary, any of the presented scenarios can really be exclude. However, a MHD wind heated by ambipolar diffusion  or a photoevaporative wind produced in an irradiated disc surface are promising scenarios. A more detailed kinematics of the H$_{2}$ cavity, within a more completed study of the connection between the low [OI] velocity component and the H$_{2}$ emission, at different positions along the jet, is desirable to better constrain these scenarios. In addition, more model predictions, and synthetic PV diagrams and maps in particular, are necessary to go further in our knowledge of this molecular component in T Tauri stars and young stellar objects in general.  %{\\bf We caution that our conclusions on shock scenarios were obtained assuming negligible turbulent mixing between the shocked wind and shocked ambient gas, following the results of \\citet{Delamarter2000}. However, magnetic fields may alter mixing in unknown ways. Fast mixing would strongly modify the predicted H$_2$ kinematics and could establish a uniform temperature of 2000~K without the need for a fine-tuned shock speed of 10 \\kms. A detailed investigation of this possibility, including the predicted H$_2$ layer thickness, is beyond the scope of the present paper but would be highly desirable, following eg. the lines of \\citet{Matsuyama2009}.}  %ADD figure of 2D uesi correction IN APPENDIX !  %\\end{itemize}  %______________________________________________________________"
    },
    "1402/1402.2992_arXiv.txt": {
        "abstract": "We present results from the coordinated broadband X-ray observations of the extreme ultraluminous X-ray source \\hoix\\ performed by \\nustar, \\xmm\\ and \\suzaku\\ in late 2012. These observations provide the first high-quality spectra of \\hoix\\ above 10\\,\\kev\\ to date, extending the X-ray coverage of this remarkable source up to $\\sim$30\\,\\kev. Broadband observations were undertaken at two epochs, between which \\hoix\\ exhibited both flux and strong spectral variability, increasing in luminosity from $L_{\\rm X} = (1.90 \\pm 0.03) \\times 10^{40}$\\,\\ergps\\ to $L_{\\rm X} = (3.35 \\pm 0.03) \\times 10^{40}$\\,\\ergps. Neither epoch exhibits a spectrum consistent with emission from the standard low/hard accretion state seen in Galactic black hole binaries, that would have been expected if \\hoix\\ harbors a truly massive black hole accreting at substantially sub-Eddington accretion rates. The \\nustar\\ data confirm that the curvature observed previously in the 3--10\\,\\kev\\ bandpass does represent a true spectral cutoff. During each epoch, the spectrum appears to be dominated by two optically thick thermal components, likely associated with an accretion disk. The spectrum also shows some evidence for a non-thermal tail at the highest energies, which may further support this scenario. The available data allow for either of the two thermal components to dominate the spectral evolution, although both scenarios require highly non-standard behavior for thermal accretion disk emission. ",
        "introduction": "Ultraluminous X-ray Sources (ULXs) are off-nuclear point sources with X-ray luminosities $L_{\\rm X} > 10^{39}$\\,\\ergps, in excess of the Eddington limit for the typical 10\\,\\msun\\ stellar-remnant black holes observed in Galactic black hole binaries (BHBs; \\eg \\citealt{Orosz03}). Multi-wavelength observations have largely excluded strong anisotropic emission as a means of substantially skewing luminosity estimates (\\eg \\citealt{Berghea10, Moon11}), thus these high luminosities require either the presence of larger black holes, either stellar remnant black holes more massive than observed in our own Galaxy (\\eg \\cite{Zampieri09}) or perhaps even the long postulated `intermediate mass' black holes (IMBHs: $10^{2}$ $\\lesssim$ \\mbh\\ $\\lesssim$ $10^{5}$\\,\\msun, \\eg \\citealt{Miller04, Strohmayer09a}), or exotic super-Eddington modes of accretion (\\eg \\citealt{Poutanen07, Finke07}). The majority of ULXs only radiate marginally in excess of $10^{39}$\\,\\ergps\\ (\\citealt{WaltonULXcat, Swartz11}), and are likely to simply represent a high luminosity extension of the stellar mass BHB population (\\citealt{Middleton13nat, Liu13nat}). Of particular interest are the rare population of extreme ULXs with X-ray luminosities $L_{\\rm X} > 10^{40}$\\,\\ergps\\ (\\eg \\citealt{Farrell09, WaltonULXcat, Jonker12, Sutton12, Walton13culx}). The extreme luminosities displayed by these sources make them the best candidates for hosting more massive black holes. For recent reviews on the observational status and the potential nature of ULXs see \\cite{Roberts07rev} and \\cite{Feng11rev}.  Previous studies have established that the 0.3--10.0\\,\\kev\\ X-ray spectra of these extreme ULXs typically show evidence for two separate continuum components (\\eg \\citealt{Miller03, Vierdayanti10, Miller13ulx}), one dominating at softer ($\\lesssim$2\\,\\kev) and the other at harder X-rays ($\\gtrsim$2\\,\\kev), potentially analogous to the disk--corona accretion components inferred for sub-Eddington BHBs (see \\citealt{Remillard06rev} for a review). However, studies focusing on the highest signal-to-noise data have found that the harder component generally shows evidence of subtle curvature in the $\\sim$3--10\\,\\kev\\ bandpass (\\citealt{Stobbart06, Gladstone09, Walton4517}), which is not observed in the coronal emission of standard sub-Eddington accretion states. A number of interpretations have since been proposed for this spectral structure, which can broadly be grouped into those that invoke thermal processes for the harder component (\\eg emission from a hot accretion disk; \\citealt{Gladstone09, Middleton11b, Sutton13uls}), which generally invoke super-Eddington accretion, and those that invoke non-thermal processes (\\eg a combination of a powerlaw continuum and relativistic reflection from the inner disk; \\citealt{Caball10}), which may still involve intermediate mass black holes (IMBHs; \\mbh\\ $\\sim 10^{2-5}$\\,\\msun). As demonstrated in \\cite{Walton4517}, these different model families predict substantially different spectra above 10\\,\\kev, in the bandpass only readily accessible with the focusing hard X-ray telescopes aboard the recently launched Nuclear Spectroscopic Telescope Array (\\nustar; \\citealt{NUSTAR}).  \\hoix\\ is one of the best studied extreme ULXs, which although known to vary in flux (\\eg \\citealt{Kong10, Vierdayanti10}) is one of the few sources (within $\\sim$5\\,Mpc) to persistently radiate at $L_{\\rm X} > 10^{40}$\\,\\ergps. Early \\xmm\\ observations revealed the possible presence of a very cool accretion disk (\\citealt{Miller03}), which may evolve in a fashion similar to the $L \\propto T^{4}$ relation expected for simple blackbody radiation (\\citealt{Miller13ulx}), and indicate the presence of a massive black hole. However, as with other extreme ULXs for which high quality data are available, \\hoix\\ shows evidence for high energy spectral curvature (\\eg \\citealt{Stobbart06, Gladstone09, Walton13hoIXfeK}). \\hoix\\ is also one of the brightest ULXs in the iron \\ka\\ bandpass, and sensitive searches have been made for absorption features that would be indicative of the massive outflows ubiquitously predicted by simulations of super-Eddington accretion (\\eg \\citealt{Ohsuga11, Dotan11}). No features are detected, with limits that require that any undetected features to be weaker than the iron absorption features resulting from the outflows in a number of sub-Eddington Galactic BHBs (\\citealt{Walton12ulxFeK, Walton13hoIXfeK}).  Reconciling these various results into a coherent picture regarding the nature of \\hoix\\ remains challenging. In order to shed further light onto the nature of the accretion in this source, we undertook a series of observations with \\nustar\\ in order to determine the nature of the high energy emission. These were coordinated with \\suzaku\\ and/or \\xmm, providing the first high quality broadband ($\\sim$0.3--30.0\\,\\kev) X-ray spectra of \\hoix. The paper is structured as follows: section \\ref{sec_red} describes our data reduction procedure, and sections \\ref{sec_spec} and \\ref{sec_var} describe the analysis performed. We discuss our results in section  \\ref{sec_dis} and summarize our conclusions in section \\ref{sec_conc}. ",
        "conclusions": "\\label{sec_conc}  We have presented results from the coordinated broadband X-ray observations of the extreme ULX \\hoix\\ performed by \\nustar, \\xmm\\ and \\suzaku\\ in late 2012. The \\nustar\\ detections provide the first high-quality spectra of \\hoix\\ at hard X-rays to date, extending our X-ray coverage up to $\\sim$30--35\\,\\kev. Observations were undertaken during two epochs, between which \\hoix\\ exhibited strong spectral variability. Neither epoch is consistent with emission from the standard low/hard accretion state, as would have been expected if \\hoix\\ harbors an IMBH accreting at substantially sub-Eddington rates; the \\nustar\\ data confirm in each case that the curvature observed previously in the 3--10\\,\\kev\\ bandpass is a true spectral cutoff. Instead, the spectrum appears to be dominated by two optically thick thermal components, likely associated with a distorted accretion disk, with a non-thermal tail also potentially detected at the highest energies probed. The data allow for either of the two thermal components to dominate the spectral evolution, although we find that both scenarios require highly non-standard behavior for behavior for thermal accretion disk emission. Further broadband observations covering a broader range of fluxes will be required to determine which component is truly dominating the observed evolution."
    },
    "1402/1402.1169_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small The study of exoplanetary atmospheres is one of the most exciting and dynamic frontiers in astronomy. Over the past two decades ongoing surveys have revealed an astonishing diversity in the planetary masses, radii, temperatures, orbital parameters, and host stellar properties of exoplanetary systems. We are now moving into an era where we can begin to address fundamental questions concerning the diversity of exoplanetary compositions, atmospheric and interior processes, and formation histories, just as have been pursued for solar system planets over the past century. Exoplanetary atmospheres provide a direct means to address these questions via their observable spectral signatures. In the last decade, and particularly in the last five years, tremendous progress has been made in detecting atmospheric signatures of exoplanets through photometric and spectroscopic methods using a variety of space-borne and/or ground-based observational facilities. These observations are beginning to provide important constraints on a wide gamut of atmospheric properties, including pressure-temperature profiles, chemical compositions, energy circulation, presence of clouds, and non-equilibrium processes. The latest studies are also beginning to connect the inferred chemical compositions to exoplanetary formation conditions. In the present chapter, we review the most recent developments in the area of exoplanetary atmospheres. Our review covers advances in both observations and theory of exoplanetary atmospheres, and spans a broad range of exoplanet types (gas giants, ice giants, and super-Earths) and detection methods (transiting planets, direct imaging, and radial velocity). A number of upcoming planet-finding surveys will focus on detecting exoplanets orbiting nearby bright stars, which are the best targets for detailed atmospheric characterization.  We close with a discussion of the bright prospects for future studies of exoplanetary atmospheres. \\\\~\\\\~\\\\~} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.6864_arXiv.txt": {
        "abstract": "We apply the dynamical systems approach to ever-expanding Bianchi type~VIII cosmologies filled with a tilted~$\\gamma$-fluid undergoing velocity diffusion on a scalar field. We determine the future attractors and investigate the late-time behaviour of the models. We find that at late times the normalized energy density~$\\Omega$ tends to zero, while the scalar potential~$\\Phi$ approaches~1 and dominates the evolution. Moreover, we demonstrate that in presence of diffusion fluids with~$\\gamma<3/2$, which includes physically important cases of dust~${(\\gamma=1)}$ and radiation~${(\\gamma=4/3)}$, are asymptotically non-tilted; the velocity of the fluid with~$\\gamma=3/2$ tends to a constant value~$0<\\bar{V}<1$; and stiffer fluids evolve towards a state of extreme tilt. Finally, we show that diffusion significantly reduces the decay rates of energy density for dust and fluids stiffer than dust~$(\\gamma \\geq 1)$; for example, at~$\\gamma=4/3$~(radiation) we obtain~${\\rho/H^2 \\propto e^{-3H_0 t}}$ at late times, while~$\\rho/H^2 \\propto e^{-4H_0 t}$ when diffusion is absent. ",
        "introduction": "The aim of theoretical cosmology is to investigate possible physical and mathematical properties of the universe. Anisotropic cosmological models, including the so-called Bianchi cosmologies \\cite{Ellis2012,Groen2007,Christiansen2008} are of great theoretical importance and have been an object of close studies; see e.\\,g.~\\cite{Hewitt2001,Coley2008,Coley2005,Hervik2005,Collins1979,Hewitt1992, Harnett1996,Hervik2004,Coley2004,Hervik2006a}, where these models have been considered using a dynamical systems approach~\\cite{Wainwright1997}. Special attention has been paid to determining the future asymptotic states of Bianchi cosmologies and describing the late-time dynamics of these universes. \\par Type VIII universes are the most general ever-expanding Bianchi cosmologies and are therefore of special interest. These models have been studied before~\\cite{Barrow2001,Ringstrom2000,Ringstrom2001,Ringstrom2003,Horwood2003}; in particular, the late-time dynamics of the model with a tilted $\\gamma$-fluid~\\cite{King1973} has been investigated in~\\cite{Hervik2006}. \\par Models where a cosmological fluid is involved into diffusional interactions with a scalar field background can be considered as a generalization of models with a cosmological constant~\\cite{Calogero2011,Calogero2012}. It has been shown, however, that models with diffusion can possess essentially different late-time dynamical properties~\\cite{Shogin2013,Shogin2014}. Also, diffusion causes some interesting effects, which have been pointed out and studied numerically in FRW~\\cite{Calogero2012,Calogero2013}, plane symmetric~$G_2$~\\cite{Shogin2013} and Bianchi class~B~\\cite{Shogin2014} universes. \\par In the current work we make a combined analytical and numerical analysis with a clear focus on the former. The aim of the paper is to confirm and generalize the results for Bianchi cosmologies obtained in~\\cite{Shogin2014} and derive the expressions describing the asymptotic dynamical features of the universe in the general case of an arbitrary fully tilted~$\\gamma$-fluid. \\par The paper is organized as follows. The evolution equations and constraints governing the dynamics of the model are given in section~\\ref{Sec:Equations}. In section~\\ref{Sec:Pre} we briefly describe the future asymptotic states of the model and also introduce some auxiliary variables. The obtained results are presented in sections~\\ref{Sec:ResTilted} and~\\ref{Sec:ResNonTilted} (the tilted and non-tilted case, respectively). Asymptotic dynamics of the geometric variables is discussed in section~\\ref{Sec:ResGeom}. Summary is presented in section~\\ref{Sec:Summary}. ",
        "conclusions": "\\label{Sec:Summary} The performed investigation shows that although cosmologies with diffusion are in some aspects similar to those with a positive cosmological constant, the late-time behaviour of these models can be substantially different. \\par The key result is that in presence of diffusion all fluids with~$\\gamma<3/2$ are asymptotically non-tilted. This range includes the cases of largest physical importance, namely dust~($\\gamma=1$) and radiation~($\\gamma=4/3$). On the contrary, in the absence of diffusion only fluids less stiff than radiation isotropize; radiation itself remains tilted in far future. \\par In cases when the fluid asymptotically becomes non-tilted the tilt decreases exponentially, both with and without diffusion. However, the decay rates are equal only for dust and fluids less stiff than dust~$(\\gamma\\leq 1).$ The velocity of fluids stiffer than dust decreases faster when diffusion is present; at late times \\begin{equation} V \\propto \\left\\{ \\begin{array}{lll} e^{-(4-3\\gamma)\\tau}, & 1<\\gamma<4/3, & \\text{without diffusion};\\\\ e^{-\\left( \\frac{3}{\\gamma}-2 \\right)\\tau}, & 1<\\gamma<3/2, & \\text{with diffusion}. \\end{array} \\right. \\end{equation} \\par Both in diffusive and diffusionless case the universes evolve towards a vacuum state. However, in presence of diffusion the asymptotic decay rates of energy density are substantially reduced for fluids stiffer than dust: \\begin{equation} \\Omega \\propto \\left\\{ \\begin{array}{lll} e^{-3\\gamma \\tau}, & 1\\leq \\gamma\\leq 4/3, & \\text{without diffusion};\\\\ e^{-3\\tau},\t\t\t& 1<\\gamma \\leq 3/2, & \\text{with diffusion}; \\end{array} \\right. \\end{equation} Dust becomes a special case when diffusion is present, the late-time dynamics of its energy density being approximated by \\begin{equation} \\Omega \\approx \\hat{K}\\tau e^{-3\\tau}+\\hat{\\Omega}e^{-3\\tau}. \\end{equation} \\par Regarding the scalar potential, diffusion does not affect the leading-order terms in the asymptotic expression. However, the higher-order terms play an important role in constructing a valid approximation in the proximity of~$\\gamma=1$; the impact of diffusion on~$\\Phi$ is thus substantial in this range. \\par The diffusion term itself decays exponentially in future; namely, at later times~$K\\propto e^{-3\\tau}$ in the range~$0<\\gamma \\leq 3/2$. \\par Finally, we should make an important remark. The characteristic time scales of evolution are substantially extended for cosmologies passing close to the point of recollapsation. This happens when the diffusion term is initially large enough~\\cite{Shogin2013,Shogin2014}. This means that the asymptotic regime does not arrive at physical times and the approximations obtained in this paper cannot in practice be applied to such universes."
    },
    "1402/1402.3393_arXiv.txt": {
        "abstract": "The massive interacting close binary system \\astrobj{V729 Cyg} (OIa+O/WN9), plausibly progenitor of a Wolf-Rayet system, is studied using new observations gathered over 65 nights and earlier published data. Radial velocity and five colour light curves are analysed simultaneously. Estimated physical parameters of the components are $M_1 = 36\\pm 3 M_{\\odot}$, $M_2 = 10 \\pm 1 M_{\\odot}$, $R_1 = 27\\pm 1 R_{\\odot}$, $R_2 = 15 \\pm 0.6 R_{\\odot}$, $\\log{(L_1/L_{\\odot})} = 5.59\\pm 0.06$, and $\\log{(L_2/L_{\\odot})}  = 4.65\\pm 0.07$. We give only the formal 1$\\sigma$ scatter, but we believe systematic errors in the luminosities, of uncertain origin as discussed in the text, are likely to be much bigger. The distance of the \\astrobj{Cygnus OB2} association is estimated as $967\\pm 48$ pc by using our newly obtained parameters. ",
        "introduction": "One of the nearest ($d < 1.5$ kpc) OB associations, \\astrobj{Cygnus OB2} contains some of the most massive and luminous stars of our Galaxy. \\astrobj{Cygnus OB2} most probably belongs to \\astrobj{Cygnus-X},  one of the richest massive star forming regions. Earlier studies showed that the \\astrobj{Cygnus OB2} region contains a number of O, B, and Wolf-Rayet stars. The total mass of the association is about 30,000 M$_{\\odot}$ (Wright et al. 2012). The age of the cluster was estimated to be in the range 1-10 Myr. Massey et al. (1995) gave an age of 1-4 Myr and recently Wright et al. (2010) give its age as 3.5-5.3 Myr. The \\astrobj{Cygnus OB2} association is a good laboratory to study star formation and stellar evolution, including the most advanced stages of the evolution of massive stars. Hence, the association has been observed in $\\gamma$-ray, X-ray, UV, optical, IR, and radio wavelengths in numerous studies. Nevertheless the interstellar medium affects the observations, especially in the UV and optical part of the electromagnetic spectrum (Rauw 2011).   The distance of the association has been calculated by various methods and these show differences. Kiminki et al. (2007) calculated the visual extinction ($A_V$) as 5.4 mag and the distance modulus as 11.3 mag, by analyzing 146 OB systems. Linder et al. (2009) have calculated the distance of the cluster as 925(25) pc from a photometric analysis of the \\astrobj{Cygnus OB2-5} system. In the literature different results derived by various methods exist (see Rauw 2011). These discrepancies possibly reflect the differences between the methods used as well as the differences in the estimation of the visual extinction. In this study, we obtain a new estimate of the distance, using our new photometric data of the eclipses   \\astrobj{V729 Cyg} (\\astrobj{Cygnus OB2-5}=BD+40$^\\circ$4220, $P=6.6$ days) is a member of the association and a massive and high-temperature contact binary (HTCB), and has been discussed in many studies since it was discovered (e.g. Wilson 1948, Wilson \\& Abt 1951, Hall 1974, Bohannan \\& Conti 1976, Leung \\& Schneider 1978, Vreux 1985, Rauw et al. 1999, Linder et al. 2009). The binary system consists of two supergiant components. Recently, Kennedy et al. (2010) reported indirect evidence for the presence of a third body in \\astrobj{V729 Cyg} with an orbital period of 6.7 years.  The UBV light variations of the system were obtained by Hall (1974) and later these light variations were used by Leung \\& Schneider (1978) to model its light curve using the Wilson-Devinney code (Wilson \\& Devinney 1971, hereafter W-D code). Leung \\& Schneider (1978) then calculated the orbital and physical parameters of the system by using Bohannan \\& Conti radial velocities. They gave the masses of the components as $M_1 = 58.7\\pm 9.1 M_{\\odot}$ and $M_2 = 13.7\\pm 6.3 M_{\\odot}$.   Vreux (1985) reported a periodic change in the H$_\\alpha$ profile. Rauw et al. (1999) measured the radial velocities of the binary components and studied the emission line profile variations caused by the stellar wind interactions. They derived the mass functions for the components as $M_1 \\sin^3 i=24.6 M_{\\odot}$ and $M_2 \\sin^3 i=6.9 M_{\\odot}$ and the mass loss rate of the stellar winds as $\\dot{M_1}=5.0\\times10^{-6}M_{\\odot}$/yr and $\\dot{M_2}=5.5\\times10^{-6} M_{\\odot}$/yr. They gave the spectral types of the components as O6.5-7~Ia+Ofpe/WN9. Linder et al. (2009) investigated the light variations of the system at 5057~{\\AA} and 6051~{\\AA}. Assuming a bright spot Linder et al. modelled its LC with the Nightfall software. The authors estimated the masses of components as $M_1 = 31.9\\pm 3.2 M_{\\odot}$ and $M_2 = 9.6\\pm 1.1 M_{\\odot}$, values with which we largely agree.   In this study newly obtained 5-colour (UBVRI) light curves of the system, published UBV light curves of Hall (1974), and radial velocity curves from Rauw et al. (1999) are analysed simultaneously and the orbital, physical parameters and the distance of the association are presented. Following the observational information given in the second section, light and radial velocity models are presented in the third section and in the fourth section physical parameters of the binary system are calculated. ",
        "conclusions": "In this study we investigated the massive interacting close binary system \\astrobj{V729 Cyg} by using our new and earlier published observations. Radial velocity and light curves of the system were solved simultaneously, and the orbital and physical parameters obtained are shown in Table~\\ref{tab:lc-results} and Table~\\ref{tab:PhyPar}. Our analysis shows that the system has a moderate contact configuration.   In their studies on X-ray and radio emission, Linder et al. (2009) and Kennedy et al. (2010) suggested that there is a third body with a period of 6.7 years orbiting the system \\astrobj{V729 Cyg} from long-term VLA observations. Kennedy et al. (2010) estimated the mass function related to the third body as 3.2  $\\rm{M_{\\odot}}$. By using this mass function value and the physical parameters given in Table ~\\ref{tab:PhyPar} the mass of the third body for $i_3$=90$^o$, 70$^o$, 50$^o$, 30$^o$ is obtained respectively as 24 $\\rm{M_{\\odot}}$, 26 $\\rm{M_{\\odot}}$, 35 $\\rm{M_{\\odot}}$, 68 $\\rm{M_{\\odot}}$. We would like to emphasize that the uncertainty in these mass values is high because of the uncertainty in the mass function.   Single massive stars above an initial mass of $8M_{\\odot}$ are known to evolve towards an explosive supernova explosion (Smartt et al. 2009). Leaving a neutron star or black hole remnant. However if a massive stars is a member of binary like \\astrobj{V729 Cyg} its evolutionary path towards the final supernovae will become significantly different. The nature of the final stellar remnant may also be altered so that while a neutron star might be expected mass transfer may lead to black hole formation. Observations of systems such as \\astrobj{V729 Cyg} will allow us to refined out binary evolutionary models to increase our understanding of the possible variety of stellar lifecycles (see Eggleton 2010, Eldridge \\& Stanway 2009, Yakut \\& Eggleton 2005, Pols et al. 1995)"
    },
    "1402/1402.6325_arXiv.txt": {
        "abstract": "We study the effects of a \\updated{detailed} dust treatment on the properties and evolution of early-type galaxies containing central black holes, as determined by AGN feedback. We find that during cooling flow episodes, radiation pressure on the dust in and interior to infalling shells of cold gas can greatly impact the amount of gas able to be accreted and therefore the frequency of AGN bursts. However, the overall hydrodynamic evolution of all models, including mass budget, is relatively robust to the assumptions on dust. We find that IR re-emission from hot dust can dominate the bolometric luminosity of the galaxy during the early stages of an AGN burst, reaching values in excess of $10^{46}$ erg/s. The AGN-emitted UV is largely absorbed, but the optical depth in the IR does not exceed unity, so the radiation momentum input never exceeds $L_{\\rm BH}/c$. We constrain the viability of our models by comparing the AGN duty cycle, \\updated{broadband luminosities}, dust mass, black hole mass, and other model predictions to current observations. These constraints force us to models wherein the \\updated{dust to metals ratios are $\\simeq 1\\%$ of the Galactic value}, and only models with a dynamic dust to gas ratio are able to produce both quiescent galaxies consistent with observations and high obscured fractions during AGN ``on'' phases. During AGN outbursts, we predict that a large fraction of the FIR luminosity can be attributed to warm dust emission ($\\simeq100$ K) from dense dusty gas within $\\leq 1$ kpc reradiating the AGN UV emission. ",
        "introduction": "It is well-established that the brightest active galactic nuclei (AGNs) and the most massive black holes reside in giant elliptical galaxies \\updated{\\citep{Dunlop_2004}}. It is also accepted that the environment provided by the galaxy, particularly the abundance of hot metal-rich gas, aging stars, and heavily depleted dust, profoundly affects the supermassive black hole (hereafter SMBH) evolution, and in turn the host system is modified by radiative and mechanical feedback \\citep[][and references therein]{Begelman+Blandford+Rees_1984, Norman+Scoville_1988, Ciotti+Ostriker_1997, Ciotti+Ostriker_2012, Ostriker+Choi+Ciotti+etal_2010, Pellegrini+Ciotti+Ostriker_2012, Vogelsberger+etal_2013}. Indeed, the interplay between the evolution of the black hole and the evolution of its host galaxy has been frequently invoked in order to explain the observed correlation between the properties of the black hole and the galaxy \\citep[e.g.][]{Springel+DiMatteo+Hernquist_2005, Ostriker+Ciotti_2005, Cattaneo+Faber+Binney+etal_2009, Debuhr+Quataert+Ma_2011}. Putting the questions we seek to answer in concrete terms: if the black hole in a galaxy such as M87 were to erupt as a quasar, what would we see? What mechanisms would govern the interaction between the accreting gas of the galaxy and the erupting black hole? What sources will contribute to the observed radiation?  A key ingredient in answering each of the above questions is the presence of dust in the galaxy. The potential importance of dust can be readily seen for several reasons. First, quasars emit at or near their Eddington limit, which classically is set by the Thomson scattering opacity of $0.4\\ {\\rm cm}^2/{\\rm g}$. By comparison, the opacity of dust grains to UV photons exceeds 1000 times this value, so the radiative forces on dust can easily overwhelm other dynamical drivers during outbursts. Second, it is widely accepted that a significant fraction of quasars are obscured. However, the giant elliptical galaxies which host AGN typically have very small optical depths, implying little dust. Therefore, not only does the dust govern the observed radiation from the galaxy, but the dust abundance itself may be a dynamic quantity that evolves in parallel with quasar ``on'' and ``off'' phases. In this work, we seek to understand the power of dust to influence the nature, evolution, and observational signatures of giant elliptical galaxies by introducing the processes that create and destroy dust within galaxies into our simulations.  One of the primary links between an AGN and its host galaxy is the immense radiative output of the AGN. The radiative feedback mechanisms between the central black hole and the accreting gas greatly influence the gas dynamics and thus have been the focus of much study \\citep[][hereafter CO07]{Binney+Tabor_1995, Ciotti+Ostriker_1997, Thompson+Quataert+Murray_2005, Ciotti+Ostriker_2007}. Because dust grains efficiently absorb UV radiation and re-radiate in the IR, the presence of grains can dramatically affect the gas dynamics \\citep[][and references therein]{Siebenmorgen+Heymann_2011}. In particular, infalling shells of cold gas (a common feature at the onset of ISM cooling episodes and detected in giant ellipticals \\citep{Werner+etal_2013}) can be supported by radiation pressure, slowing accretion, and altering the subsequent evolution. In turn, this evolution is related to the so-called ``positive feedback'' mode in which feedback from the AGN enhances star formation \\citep[CO07,][and references therein]{Nayakshin+Zubovas_2012}. Observations have hinted at the presence of radiation-supported cold gas around AGN \\citep{Vasudevan+etal_2013} though in the majority ($\\simeq95$\\%) of cases the gas is below its effective Eddington limit \\citep{Raimundo+etal_2010}. Nevertheless, the latter authors note the inevitability of radiation-supported gas and its possible role in AGN feedback. Radiation pressure on dust grains may be instrumental in driving galactic winds \\citep{Coker+Thompson+Martini_2013}, though there is disagreement as to whether this could be a substantial effect \\citep{Socrates+Sironi_2013}.  In a recent series of papers \\citep{Ciotti+Ostriker+Proga_2009, Shin+Ostriker+Ciotti_2010, Ciotti+Ostriker+Proga_2010}, 1D hydrodynamical simulations were used to study the mechanisms of AGN feedback. The dust physics was implemented in a very simple, phenomenological way. As sputtering effectively destroys dust in regions of hot gas \\citep{Draine+Salpeter_1979}, these previous papers approximated the sputtering by reducing the dust to gas ratio in hot gas by roughly two orders of magnitude relative to Galactic values in accord with observations of dust in elliptical galaxies. In practice, a factor inversely proportional to the gas temperature multiplies the fiducial absorption coefficient in the UV, optical, and IR.  In this paper, we improve significantly this aspect of the input physics. We seek to clarify the role of dust grains in AGN feedback by developing and numerically implementing the basic equations for dust production and destruction. \\updated{While a similar but simpler version of the dust physics described here was implemented in the 2D simulations of \\citet{Novak+Ostriker+Ciotti_2012}, we present a more generalized treatment and focus principally on the effects of different models of dust abundance on the hydrodynamical evolution of the simulation, both in its galaxy-scale properties and accretion physics. Additionally, we discuss the ability or inability of models to reproduce the observational properties of giant elliptical galaxies in the infrared during both quiescent and accretion phases.}  For the present hydrodynamical simulations we use an updated version of the 1D code used in our previous studies, with the relevant modifications detailed in Section~\\ref{sec:simulations}. In particular, a major upgrade has been made in the solution of the radiative transfer equation. We note that in a recent related paper \\citep{Novak+Ostriker+Ciotti_2012} the new ``two-streams'' approach is also tested in the context of 2D hydrodynamical simulations.  We restrict our focus to the secular evolution of the galaxy. While processes such as major and minor mergers or inflows of cold gas from the IGM certainly occur and can account for a variety of observations, we are most interested in which aspects of the evolution of these galaxies can be accounted for by purely secular processes, i.e. driven by well-understood stellar evolutionary processes. We find that the AGN duty cycle, its time dependence, the black hole mass, and the IR luminosity are all naturally explained with our purely secular model.  The paper is organized as follows: in Section~\\ref{sec:physics} we describe the dust grain physics employed in our simulations as well as a discussion of two-stream radiative transfer. In Section~\\ref{sec:simulations} we describe the suite of simulations conducted under different models for the grain abundance. In Section~\\ref{sec:survey}, we summarize the results of the simulations and compare against observations to select the most viable models. We also discuss the effects of dust on the properties of AGN outbursts. We then discuss the most viable models in the context of observations in Section~\\ref{sec:observations}, focusing on both observational signatures predicted by our model as well as how current observations constrain the survival, growth, and radiative properties of dust in elliptical galaxies with AGN. We summarize the implications of these comparisons in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} By implementing a more physically-based dust treatment into 1D hydrodynamical simulations of the evolution of massive elliptical galaxies, we are able to link the computed IR emission from the galaxy during various stages of secular evolution with observations of IR emission. These models are capable of attaining LIRG and ULIRG-like phases of high IR emission without needing to invoke non-secular processes.  Despite the differing assumptions on dust abundance, the simulated galaxies illustrated a remarkably robust mass budget-- in each simulation, the vast majority of the gas in the galaxy was expelled in outflows, about $10\\%$ was turned into stars, a few percent was accreted onto the central black hole, and a few percent remained as gas. The black hole growth is consistent both with current determinations of the Magorrian relation and the empirical fact that quasar ``on'' periods decline in frequency with decreasing redshift.  Our most physical dust models are able to reconcile the low observed dust abundance of quiescent galaxies (dust to gas ratios of $\\simeq10^{-4}$) with presence of heavily obscured quasars through grain growth and reduced sputtering rates in cold gas. Additionally, optically-thick gas was able to oscillate between accretion and outflow phases for tens of Myr, resulting in sustained periods of large IR luminosity consistent with LIRGs and ULIRGs. However, at variance with \\citet{Debuhr+Quataert+Ma_2011}, $\\tau_{\\rm IR}$ never exceeds unity and the momentum imparted to the dust gas never exceeds $L_{\\rm BH}/c$.  We identify two distinct types of AGN bursts common to all models-- short-duration optically thin bursts that eventually culminate to a single large, complex burst that is largely optically thick. A clear prediction of this work is the presence of infrared emission from $\\simeq100$ K dust grains in the inner $\\simeq1$ kpc of massive galaxies during AGN bursts.  The presence of dust grains in accreting gas was also found to impact the star formation processes in the galaxy-- AGN feedback and consequent quenching of star formation was enhanced in models with more dust. Similarly, dusty models also accrete less gas and have shorter duration bursts. Irrespective of our dust treatment, we find periods of ``positive feedback'' on star formation in which AGN activity precipitates a brief period of active star formation.  An inherent limitation of 1D simulations is the inability to account for fragmentation of gas. The influence of dust in this case, particularly in its role of preventing gas from accreting, has yet to be determined using a detailed physical prescription for the dust abundance. The formalism laid out in this work can be easily generalized to higher dimensional simulations, and given the importance of dust not only in the dynamics but also observational signatures, doing so may shed additional light on evolution of these galaxies."
    },
    "1402/1402.0053_arXiv.txt": {
        "abstract": "Efficient parameter estimation is critical for Gravitational-Wave astronomy. In the case of compact binary coalescence, the high dimensional parameter space demands efficient sampling techniques---such as Markov chain Monte Carlo (MCMC). A number of degeneracies effectively reduce the dimensionality of the parameter space and, when known, can render sampling algorithms more efficient with problem-specific improvements.  We present in this paper an analytical description of a degeneracy involving the extrinsic parameters of a compact binary coalescence gravitational-wave signal, when data from a three detector network (such as Advanced LIGO/Virgo) is available. We use this new formula to construct a jump proposal, a framework for a generic sampler to take advantage of the degeneracy. We show the gain in efficiency for a MCMC sampler in the analysis of the gravitational-wave signal from a compact binary coalescence. ",
        "introduction": "Among the sources of gravitational waves (GWs), inspiralling binary systems of compact objects, neutron stars (NSs) and/or black holes (BHs) in the mass range $\\sim 1\\,\\Ms - 100\\,\\Ms$ stand out as likely to be detected and relatively easy to model.  For the network of ground-based laser interferometers \\cite{Cutler:2001}, LIGO (Laser Interferometer Gravitational-wave Observatory) \\cite{2009NJPh...11g3032A} and Virgo \\cite{Acernese:2008b}, currently undergoing upgrades, the detection-rate estimates for compact object binaries, although uncertain, are expected to be about $70\\, \\mathrm{yr}^{-1}$ \\cite{ratesdoc}.  The detection of a gravitational-wave event is challenging and will be a rewarding achievement by itself. After such a detection, measurement of source properties holds major promise for improving our astrophysical understanding of these sources and requires efficient methods for parameter estimation.  This is a complicated problem because of the large number of parameters ($15$ for spinning compact objects in a quasi-circular orbit) and the quasi-degeneracies between them \\cite{Raymond:2009}, the significant amount of structure in the parameter space, and the particularities of the detector noise.  We analyse the signal produced during the inspiral phase of two compact objects of masses $M_{1,2}$ in quasi-circular orbit.  A circular binary inspiral with both compact objects spinning is described by a 15-dimensional parameter vector $\\vec{\\lambda}$. A possible choice of independent parameters with respect to a fixed geocentric coordinate system is: \\be \\vec{\\lambda} = \\{m_1,m_2,d,t_c,\\phi,\\alpha,\\delta,\\iota,\\psi, a_\\mathrm{spin1},\\theta_\\mathrm{spin1},\\phi_\\mathrm{spin1},a_\\mathrm{spin2},\\theta_\\mathrm{spin2},\\phi_\\mathrm{spin2}\\} \\label{e:lambda} \\ee where $m_1$ and $m_2$ are the masses of the heaviest and lightest members of the binary, respectively; $d$ is the luminosity distance to the source; $\\phi$ is an integration constant that specifies the gravitational-wave phase at a reference frequency; the time of coalescence $t_c$ is defined with respect to the centre of the Earth; $\\alpha$ (right ascension) and $\\delta$ (declination) identify the source position in the sky; $\\iota$ defines the inclination of the binary with respect to the line of sight; and $\\psi$ is the polarisation angle of the waveform.  The spins are specified by $0 \\le a_\\mathrm{spin_{1,2}} \\equiv S_{1,2}/M_{1,2}^2 \\le 1$ as the dimensionless spin magnitude, and the angles $\\theta_\\mathrm{spin1,2}$, $\\phi_\\mathrm{spin1,2}$ for their orientations with respect to the line-of-sight.  It is convenient to define two families of parameters. The intrinsic parameters: \\be \\overrightarrow{\\lambda}_{intrinsic} = \\{m_1,m_2, a_\\mathrm{spin1},\\theta_\\mathrm{spin1},\\phi_\\mathrm{spin1},a_\\mathrm{spin2},\\theta_\\mathrm{spin2},\\phi_\\mathrm{spin2}\\} \\ee are required for the computation of the gravitational wave in any reference frame. The extrinsic parameters: \\be \\overrightarrow{\\lambda}_{extrinsic} = \\{d,t_c,\\phi,\\alpha,\\delta,\\iota,\\psi\\} \\ee control the projection of the gravitational wave onto the geocentric reference frame, in which we can compute the response of each detector with \\eqnref{eq:beam_pattern}.  Given a network comprising $n_\\mathrm{det}$ detectors, we assume that the data collected at the $i-$th instrument ($i = 1,\\dots, n_\\mathrm{det}$) is given by $x_i(t) = n_i(t) + h_i(t;\\vec{\\lambda})$, where $h_i(t;\\vec{\\lambda})$ is the gravitational-wave signal (see \\eqnref{eqn:signal}), and $n_i(t)$ is the detector noise (here assumed to be stationary and normally-distributed).  The equations governing the response of an observatory to gravitational waves have long been known, see for instance \\cite{1973grav.book.....M} and references therein. To illustrate the degeneracy present in this response we use Markov chain Monte Carlo (MCMC) methods to determine the multi-dimensional \\emph{posterior} probability-density function (PDF) of the unknown parameter vector $\\vec{\\lambda}$ in equation~\\ref{e:lambda}, given the data sets $x_i$ collected by a network of $n_\\mathrm{det}$ detectors, a model $M$ of the waveform and the \\emph{prior} $p(\\vec{\\lambda})$ on the parameters. One can compute the probability density via Bayes' theorem \\be p(\\vec{\\lambda}|x_j,M) = \\frac{p(\\vec{\\lambda}|M) \\, p(x_j|\\vec{\\lambda},M)}{p(x_j|M)}\\,, \\label{e:jointPDF} \\ee where \\be L \\equiv p(x_j|\\vec{\\lambda},M) \\propto \\exp\\left( <x_j|h_j(\\vec{\\lambda})>-\\frac{1}{2}<h_j(\\vec{\\lambda})|h_j(\\vec{\\lambda})> \\right) \\label{e:La} \\ee is the \\emph{likelihood function}, which measures the probability (under the noise distribution) of getting data $x_j$ given a signal $h_j$. The term $p(x_j|M)$ is the \\emph{marginal likelihood} or \\emph{evidence}. In the previous equation \\be <x|y>=4Re\\left( \\int_{f_{\\rm low}}^{f_{\\rm high}}\\frac{\\tilde{x}(f)\\tilde{y}^{*}(f)}{S_j(f)}\\,\\mathrm{d}f \\right) \\label{e:prod} \\ee is the \\emph{overlap} of signals $x$ and $y$, $\\tilde x(f)$ is the Fourier transform of $x(t)$, and $S_j(f)$ is the noise power-spectral density in detector $j$. The likelihood computed for the injection parameters $\\mathcal{L}_\\mathrm{inj}=p(x_j|\\vec{\\lambda}_\\mathrm{inj},M)$ is then a random variable that depends on the particular noise realisation $n_j$ in the data $x_j=h(\\vec{\\lambda}_\\mathrm{inj})+n_j$. The injection parameters are the parameters of the waveform template added to the noise.  To combine observations from a network of detectors with uncorrelated noise realisations we have the likelihood $p(\\vec{x}|\\vec{\\lambda},M) = \\prod_{a=1}^{n_\\mathrm{det}}\\, p(x_j|\\vec{\\lambda},M)\\,$, for $\\vec{x} \\equiv \\{x_j: j = 1,\\dots,n_\\mathrm{det}\\}$ and \\be p(\\vec{\\lambda}|\\vec{x},M) = \\frac{p(\\vec{\\lambda}|M)\\, p(\\vec{x}|\\vec{\\lambda},M)}{p(\\vec{x}|M)}. \\label{e:Bayes} \\ee  The numerical computation of the PDF involves the evaluation of a large, multi-modal, multi-dimensional integral. MCMC methods \\citeaffixed[and references therein]{gilks_etal_1996,GelmanCarlinSternRubin}{\\emph{e.g.}} have proved to be especially effective in tackling this numerical problem.  In the Markov chain Monte Carlo algorithm, a Markov chain crawls around the parameter space according to a specific set of rules: \\begin{itemize} \\item{}At iteration n, the chain is in the state $\\vec{\\lambda}_n$. Choose a proposal state $\\vec{\\lambda}_k$ with probability $p(\\vec{\\lambda}_k|\\vec{\\lambda}_n)$. \\item{}Compute the acceptance probability $p_a$: \\be p_a=\\min\\left\\{ 1, \\frac{p(\\vec{\\lambda}_k\\,|\\vec{x},M)p(\\vec{\\lambda}_n|\\vec{\\lambda}_k)}{p(\\vec{\\lambda}_{n}|\\vec{x},M)\\,p(\\vec{\\lambda}_k|\\vec{\\lambda}_n)} \\right \\} % \\ee \\item{}Accept $\\vec{\\lambda}_k=\\vec{\\lambda}_{n+1}$ as the new state of the chain with probability $p_a$, otherwise $\\vec{\\lambda}_{n+1}=\\vec{\\lambda}_{n}$ (with probability $1-p_a$) \\end{itemize}  The distribution of parameters in the set of states $\\left\\{\\vec{\\lambda}_n\\right\\}$ of the chain following this procedure converges towards the posterior distribution as $n \\to \\infty$. Note that for any proposal to be included in this algorithm, the ratio \\be r\\equiv\\frac{p(\\vec{\\lambda}_n |\\vec{\\lambda}_k)}{p(\\vec{\\lambda}_k | \\vec{\\lambda}_n)}, \\ee needs to be computed, see \\secref{sec:detail_balance}.    We derive for the first time in the literature a proposal that generates jumps in parameter space that exploit a near-degeneracy in the detector responses for the three-detector case.  Using such a proposal in the context of an MCMC generates moves that efficiently explore the extrinsic dimensions of the posterior distribution function, even when the posterior is multi-modal with widely separated, narrow peaks in the extrinsic dimensions.  In this paper we first present the existing degeneracies involving the extrinsic parameters describing a binary coalescence in \\secref{sec:deg}. In \\secref{sec:eqn} we present the equations which we solve in \\secref{sec:sol} to generate proposed moves. In \\secref{sec:detail_balance} we address detailed balance. We apply our proposal in our Markov chain Monte Carlo algorithm and describe the results in \\secref{sec:res}. Finally we conclude in \\secref{sec:end}. ",
        "conclusions": "\\label{sec:end}  We described in this paper a proposal which allows for a much better exploration of the extrinsic parameter space for non-spinning gravitational wave signals. It should still be helpful in the spinning case, whose leading-order behavior mirrors the non-spinning case; we plan to test this in future work.  It may be possible that using an approximation beyond Quadrupole instead of \\eqnref{eqn:const} leads to a better handle on the spinning case where there is not simple relation between $H_+$ and $H_{\\times}$.  It may also be necessary to include some intrinsic parameters to construct a more efficient proposal for spinning analyses, as precession of the orbital plane couples the spin parameters to the inclination."
    },
    "1402/1402.6169_arXiv.txt": {
        "abstract": "In this work we test the most widely used methods for fitting the composition fraction in data, namely maximum likelihood, $\\chi^2$, mean value of the distributions and mean value of the posterior probability function.  We discuss the discrimination power of the four methods in different scenarios: signal to noise discrimination; two signals; and distributions of Xmax for mixed primary mass composition. We introduce a \"distance\" parameter, which can be used to estimate, as a rule of thumb, the precision of the discrimination. Finally, we conclude that the most reliable methods in all the studied scenarios are the maximum likelihood and the mean value of the posterior probability function. ",
        "introduction": "One of the main challenges of the astroparticle physics reasearch is to determine the chemical composition of the cosmic rays that reach the earth. At the highest cosmic ray energies, analysis of composition must be made with indirect methods by using the evolution of cosmic ray showers in the atmosphere.  In this note, we use four different methods to estimate the composition given a single variable (e.g. the $X_{\\rm max}$ distribution) by using Bayesian methods. A similar work, using Monte Carlo techniques, to study the efficiency of different discriminators can be found in \\cite{Durso}.  The note is as follows: in section \\ref{Methods} we discuss the four different methods and we apply them in simple analytical cases in section \\ref{Analytical}. In section \\ref{Examples}, we apply the different methods to two physical cases Signal to Noise discrimination and two overlapped signals. Finally, we evaluate the methods for the specific example of $X_{\\rm max}$ distributions with realistic probability densities and we get a possible composition of the cosmic rays in section \\ref{Xmax}.  To be concrete, we will consider a two composition scenario. Although the methods discussed can be easily generalized to include more than two distributions. ",
        "conclusions": "\\label{end}  We have studied different estimators for the evaluation of composition fraction between two distributions.  We have shown that the best methods are the maximum and the mean of the probability distribution.  The $\\chi^2$ method gives results comparable to the maximum likelihood estimator, if the number of events is large, but by rebinning the results can be misleading. Also, we obtained the remarkable results that with few events, the mean value of the probability distribution is the best estimator.  We found a measure of distance between the two probabilities which gives us an estimation of the discrimination power for two distributions. If the distance $d_1$ is small, the discrimination between the two compositions will be poor. If the distance is large it will be optimal. We have shown that as a ``rule of thumb'' the discrimination power scales as $1/\\sqrt{d_1 N}$ with $N$ the number of events. Generalization of these methods to include an arbitrary number of components is straightforward and will be discussed separately."
    },
    "1402/1402.3861_arXiv.txt": {
        "abstract": "The present paper highlights the data analysis status of the OPERA experiment. The experiment was designed to perform the neutrino interactions analysis on event-by-event basis, and optimized to search for $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ oscillation in appearance mode, also allowing to perform a $\\nu_{e}$ appearance search.  We review the data simulation and the analysis chains implemented to search for $\\nu_{\\tau}$ interactions. The main kinematical parameters sensitive to the neutrino flavor are discussed, the uncertainties of the event parameter estimation are reviewed, and the main sources of background for the $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ oscillation search are examined. The topologies of the two first $\\nu_{\\tau}$ candidate events are presented.  Finally, we review the status of the $\\nu_{e}$ appearance search and present the constraints set by the OPERA experiment on the mixing angle $\\theta_{13}$ and on the LSND/MiniBooNE anomaly. ",
        "introduction": "\\label{sec:intr} An idea of neutrino oscillations was introduced by B.M. Pontecorvo long ago \\cite{pon57}. However, it took over 40 years to get the first definite experimental result in favor of oscillations \\cite{fuk98}. Now it is firmly established that the number of the ``active'' neutrino flavors amounts to 3 (formal fit gives 2.998$\\pm$0.008 \\cite{hag02}), thus oscillations between these states are described by 3$\\times$3 Pontecorvo-Maki-Nakagawa-Sakata (PMNS) matrix \\cite{mak62}.  Further experiments with solar \\cite{ahm02}, atmospheric \\cite{amb98,all05}, reactor \\cite{ara05} and accelerator \\cite{ahn06,mih06,ada08} neutrinos gave additional support to the oscillation hypothesis. Recently the value of the last previously unknown mixing angle $\\theta_{13}$ was measured in the Daya Bay ex\\-pe\\-ri\\-ment \\cite{an12}. Some phenomenological models introduce and some experiments (e.g. \\cite{aqu01,aqu12}) support the existence of $m$ additional (``sterile'') neutrino states that may participate in oscillations; in that case the PMNS matrix is $3+m$-dimensional.  Other unsolved problems, besides the question of sterile neutrinos existence, are those of neutrino mass hierarchy and of CP violation in the leptonic sector. The decisive proof of $\\nu_{\\tau}$ appearance in $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ oscillation is also still absent; only recently the Super-Kamiokande experiment found the evidence for this effect \\cite{abe13}.  The OPERA experiment was designed to search for $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ oscillation on event-by-event basis by observing $\\tau$ decays. OPERA might be also handy in constraining the active-sterile neutrino mixing parameters, and the so-called non-standard interaction (NSI) parameters \\cite{ter11}, that may give rise to $\\nu_{\\tau}$ component in the initially pure $\\nu_{\\mu}$ beam.  In the present paper we review the data analysis status and the results of the OPERA experiment. Much more detailed description of the data analysis chain is given in \\cite{aga13a}. For the OPERA hardware reference see \\cite{acq09}. The status of emulsion scanning is presented in the other OPERA contribution in this volume \\cite{she13}.  In sec.~\\ref{sec:chain} we present a sketch of the OPERA analysis procedure. In the following sections some details of simulations and analysis are given: the parameters sensitive to the primary neutrino flavor are introduced in sec.~\\ref{sec:param}, the emulsion detector capabilities to reconstruct the neutrino interaction kinematical parameters are discussed in sec.~\\ref{sec:est}, the background studies are summarized in sec.~\\ref{sec:bckg}. In sec.~\\ref{sec:cand} we present the two first $\\nu_{\\tau}$ candidates, and in sec.~\\ref{sec:elec} --- the results concerning the $\\nu_{e}$ appearance in the OPERA experiment. ",
        "conclusions": "\\label{sec:conc} We have reviewed the present status of the OPERA experiment aimed on event-by-event $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ oscillation search in appearance mode. Three $\\nu_{\\tau}$ candidates were observed so far. The main sources of background were carefully studied. The complete simulation/analysis framework OpRelease was recently established, allowing to perform more realistic MC studies. Constraints on $\\theta_{13}$ and the LSND/MiniBooNE anomaly were set profiting of the excellent particle identification capabilities of the OPERA detector."
    },
    "1402/1402.1749_arXiv.txt": {
        "abstract": "Scenarios of the cosmological evolution  are studied by using an equation of state (EoS) having points where the specific enthalpy of the cosmological fluid vanishes. A large class of barotropic EoS's admits, depending upon initial conditions, analogues of the ``Big Rip\", as well as solutions describing exponential inflation followed by usual matter dominance; their classification is proposed. We discuss extensions to more general two-parametric EoS dealing with a pre-inflationary evolution and yielding stages with both increasing and decreasing energy density as a function of time. Possible cosmological scenarios with transitions from collapse to an expanding Universe or a closed oscillating one, without reaching a singularity are included. ",
        "introduction": "\\label{Introduction} The Standard Cold Dark Matter cosmological model with $\\Lambda\\ne 0$ ($\\Lambda$CDM model)  describes a huge amount of observational data \\cite{WMAP,Plank} that determine the cosmological fraction of the baryonic matter, the cold dark matter and the dark energy. On the other hand, the $\\Lambda$CDM model leads to the well-known problems of horizon and of spatial flatness (see, {\\it e.g.} \\cite{Weinberg,Linde}). These problems are currently resolved by introducing the inflationary period \\cite{inflation} in the early Universe, leading to a high degree of isotropy of the cosmic microwave background and spatial flatness. This can be achieved either within modifications of General Relativity or by taking into account some additional physical fields \\cite{Sahni,Linde}, which are manifest at cosmological scales.  Traditionally, the bulk of matter in the Universe is represented by a sum of baryonic matter, unknown dark matter (DM) with zero (or very small) pressure and dark energy (DE). In the hydrodynamical picture, the total pressure is a sum $p=p_B+p_{DM}+p_{DE}$. This picture can be complemented by introducing a cosmological ({\\it e.g.}, scalar) field. On the other hand, at present there are no theoretical arguments that could single out a unique scalar field potential, a Lagrangian of a modified gravitation theory etc. In this situation, qualitative considerations of cosmological scenarios often use a phenomenological equation of state (EoS) describing on the average the whole matter  in terms of the local frame energy density $e$ and other thermodynamical parameters (see \\cite{eos_models} and references therein). A widely accepted model for the dark energy is a homogeneous perfect fluid with $p=w e$. Cosmic acceleration requires that $w<-1/3$; recent WMAP and Planck results \\cite{WMAP,Plank} are consistent with the value $w \\approx -1$, however ``phantom DE\",  for which $w<-1$ \\cite{BigRip}, is not completely ruled out \\cite{phantom}.  A number of generalizations  dealingwith either explicit forms of EoS or having a parametric form exist \\cite{eos_models}. In this paper we {use an} EoS containing a contribution {from} ordinary matter with { a} linear part $p=w e$, $w=const$ ({\\it e.g.}, due to ``hot\" matter with $w=1/3$) and a nonlinear term depending on the total energy density $e$. This EoS is inspired by the well-known bag models of the quark-hadron (de)confinement phase transition when quarks and gluons coagulate to form hadrons (see, {\\it e.g.}, \\cite{Shurik,JKS,JKS1}). These models involve the so called ``bag\" pressure constant $\\bar{B}$ that can be interpreted as the effect of the quantum-chromodynamical vacuum. A generalization in which the constant $\\bar{B}$ is replaced by an increasing function of temperature $T$, was suggested by K\\\"allman \\cite{Kallman}. This modification of the EoS was re-derived and its hydrodynamical and cosmological consequences have been discussed in a number of papers (see \\cite{JKS, JKS1, ECh} and references therein). The possibility that a small fraction of colored objects escape hadronization, surviving as islands of free coloured particles, called quark nuggets was studied {\\it e.g.} in Ref.~\\cite{Odessa}.  Possible variation of the cosmological constant with temperature $T$ or energy density $e$ was discussed {\\it e.g.} in Ref. \\cite{Lambda}. Of interest is the connection of this phenomenon with the variable quark bag ``constant\" $B$ considered in the present paper. Below we extend the idea of the variable vacuum pressure $B$ to very early times of the Universe, although we do not exclude that the nonlinear part of the equation of state can have different origin. We assume that $B$ may be a function of $e$. This is applied to equations of the homogeneous isotropic Universe  (Section \\ref{hydrodynamical_model}) dealing with a barotropic EoS. In Section \\ref{section3} we discuss more general equations of state that involve two thermodynamical parameters, namely, the energy density and the specific volume. The results are summarized in Section \\ref{conclusions}. ",
        "conclusions": "We have investigated a  class of nonlinear equations of state that have points of zero enthalpy; to do so, we used general  qualitative properties known from the theory of ordinary differential equations. In the case of the barotropic EoS occurrence of such point leads to the possibility of a period of rapid expansion with almost constant energy density. Under appropriate initial conditions, this period can last as long as needed to provide the necessary size of the causally connected regions in future, ensuring, {\\it e.g.}, an adequate solution to the well-known problem of horizon etc.  Under general assumptions, we have classified possible cosmological scenarios, including a  number of those without the cosmological singularity. In case of an open and spatially flat metrics, the Universe expands either forever or undergoes a Big Rip. In case of the closed Universe there can be also transitions from collapse to expansion or vice versa. A closed oscillating Universe without reaching a singularity is also possible.  Our findings deal with  simple EoS allowing to discuss different scenarios of the early cosmological evolution. Nevertheless, if we assume that these considerations have relevance to reality, one can speculate on why our Universe starts with small deviations from the state with zero enthalpy, for example due to quantum fluctuations near this state. These fluctuations can start either ``phantom\" cosmological evolution leading to a kind of the Big Rip, or ``normal\" expansion  with a transition to the standard $\\Lambda$CDM  model described in textbooks (see, {\\it e.g.}, \\cite{Weinberg}). To have a sufficiently long period of inflation, a kind of fine tuning is required ({\\it e.g.}, approach to the point $h=0$), compelling  to invoke a sort of the anthropic principle. However, such a tuning is not more restrictive than, {\\it e.g.}, the use of  equation of state $p=w e$ with $w=-1$.  The above-mentioned requirement of the fine tuning is relaxed if we consider a two-parametric EoS. In this case, an infinity of solutions crossing states with zero enthalpy may exist; and there is an immense freedom in choosing the evolution as the scale factor $a\\to 0$. In case of a sufficiently weak dependence of $p(e,v)$ on $v$ it is natural to assume that the behavior of trajectories near points of zero enthalpy will be similar to that of Section \\ref{one-parametric} and there is a sufficiently long period of inflation.  Furthermore, one may speculate about a multi-component fluid. We know that at present some form of DE dominates, but moving back in time its fraction becomes negligible (within the framework of the $\\Lambda$CDM model). On the other hand, for very early (post-Planckian) times, the other form of  DE must have dominated to provide inflation. Within such a scenario, one may have different inflation epochs corresponding to the domination of different DE components (cf. multiple stages of inflation in Ref. \\cite{schism}). This is consonant with the idea of the existence of a ``mini-inflation'' due to the existence of metastable states in strongly interacting matter \\cite{JKS,JKS1}, see also \\cite{Shurik}. However, any smooth interpolation between these different inflation epochs remains a challenge for the theory."
    },
    "1402/1402.3927_arXiv.txt": {
        "abstract": "{We present a Monte Carlo simulation for the scattering of light in the case of an isotropic light source. The scattering phase functions are studied particularly in detail to understand how they can affect the multiple light scattering in the atmosphere. We show that although aerosols are usually in lower density than molecules in the atmosphere, they can have a non-negligible effect on the atmospheric point spread function. This effect is especially expected for ground-based detectors when large aerosols are present in the atmosphere.} ",
        "introduction": "Light coming from an isotropic source is scattered and/or absorbed by molecules and/or aerosols in the atmosphere. In the case of fog or rain, the single light scattering approximation -- when scattered light cannot be dispersed again to the detector and only direct light is recorded -- is not valid anymore. Thus, the multiple light scattering -- when photons are scattered several times before being detected -- has to be taken into account in the total recorded signal. Whereas the first phenomenon reduces the amount of light arriving at the detector, the latter increases the spatial blurring of the isotropic light source. This effect is well known for light propagation in the atmosphere and has been studied by many authors, e.g.~\\cite{Kopeika_3}. The problem of light scattering in the atmosphere has no analytical solutions and Monte Carlo simulations are usually used to study light propagation in the atmosphere. A multitude of Monte Carlo simulations have been developed in the past years, all yielding similar conclusions: aerosol scattering is the main contribution to atmospheric blur, atmospheric turbulence being much less important. A significant source of atmospheric blur is especially aerosol scatter of light at near-forward angles~\\cite{Kopeika_3,Reinersman}. The multiple scattering of light is affected by the optical thickness of the atmosphere, the aerosol size distribution and the aerosol vertical profile. Whereas many works have studied the effect of the optical thickness, the aerosol blur is also very dependent on the aerosol size distribution, and especially on the corresponding asymmetry parameter of the aerosol scattering phase function. The purpose of this work is to better explain the dependence of the aerosol blur on the aerosol size, and its corresponding effect on the atmospheric point spread function. Section~\\ref{sec:modelling} is a brief introduction of some quantities concerning light scattering, before describing in detail the Monte Carlo simulation developed for this work. Then, in Section~\\ref{sec:global_view}, we explain how different atmospheric conditions affect the multiple scattering contribution to the total light arriving at detectors within a given integration time across all space. This result is finally applied to the point spread function for a ground-based detector in Section~\\ref{sec:ground_telescope}. ",
        "conclusions": "A Monte Carlo simulation for the scattering of light has been created and used to observe atmospheric aerosol effects on the percentage of indirect light collected by detectors. The study began with a general description of the dispersion of scattered photons in different atmospheric conditions. It was found that for an increased value of the asymmetry parameter $g$ ({\\it i.e.}\\ a larger aerosol size), a greater accumulation of scattered photons delayed with respect to direct photons by only a few hundreds of ns is found. The principal argument presented in this work is that, even for a low density of aerosols in the atmosphere, the ratios of indirect to direct photons detected can be comparable or greater to those caused by molecules. In particular, the value of detection time $t_{\\mathrm{det}}$ is proved to play an important role on the relative effects of molecules and aerosols on the ratio. This phenomenon is also used to estimate the aerosol size distribution, especially for very large aerosols. The technique is described in detail in~\\cite{Zacanti,Trakhovsky_1,Trakhovsky_2}.  In addition, this aerosol size effect could still solve some unsolved experimental observations as the measurement done at the Pierre Auger Observatory~\\cite{PAO,MyICRC} a few years ago. Indeed, part of the point spread function measured by ground-based telescopes is still not fully understood, {\\it i.e.}\\ cannot be reproduced in simulations~\\cite{ICRC_Julia,MS_Assis}. One of possible explanations could be an additional contribution coming from a population of large aerosols present in the atmosphere."
    },
    "1402/1402.4500_arXiv.txt": {
        "abstract": "{ In the simplest models of asymmetric dark matter (ADM) annihilation signals are not expected, since the DM is non-self-conjugate and the relic density of anti-DM is negligible. We investigate a new class of models in which a symmetric DM component, in the `low-mass' 1-10 GeV regime favoured for linking the DM and baryon asymmetries, is repopulated through decays. We find that, in models without significant velocity dependence of the annihilation cross section, observational constraints generally force these decays to be (cosmologically) slow. These late decays can give rise to gamma-ray signal morphologies differing from usual annihilation profiles. A distinctive feature of such models is that signals may be absent from dwarf spheroidal galaxies. } ",
        "introduction": "\\label{sec:intro}   Asymmetric dark matter (ADM) provides a well motivated framework for light DM and is an intriguing alternative to the usual WIMP scenario. In such models the DM, which we denote here $B'$, has an (approximately) conserved quantum number (which we also call $B'$). The relic density is determined by a particle-antiparticle asymmetry between $B'$ and $\\overline {B'}$, in direct analogy to baryons \\cite{review}. If the DM has a similar mass to the proton $m_{B'}\\sim m_{p}$ and the hidden and visible sectors are connected via portal operators which violate $B$, $L$ and $B'$, but conserve some linear combination, then this can explain the cosmological coincidence $\\Omega_{\\rm DM}\\sim5\\Omega_{B}$. In contrast, the accidental proximity of $\\Omega_{\\rm DM}$ and $\\Omega_{B}$ in conventional DM scenarios seems unreasonable given that the relic density of DM is determined by freeze-out, whereas the baryon density is set by CP-violating decays of out-of-equilibrium states and these two mechanisms are typically unrelated.  While we shall not specify the UV physics that produces these particle asymmetries, there are several mechanisms that can {\\em cogenerate} equal magnitudes for $\\eta_{B'}$, $\\eta_{B}$ and/or $\\eta_{L}$ (see e.g.~\\cite{cogenesis}). Alternatively, an asymmetry could be generated in a single quantum number and subsequently shared via processes that violate the individual global symmetries, leading to comparable asymmetries. For the case $\\eta_{B'}\\simeq\\eta_{B}$, to account for the  DM relic density  one requires the DM mass to be around 5 GeV. Inefficient sharing, or bias generation, of the DM and baryon asymmetries may easily lead to $\\mathcal{O}(1)$ deviations between $\\eta_{B'}$ and $\\eta_{B}$ and hence the DM mass could reasonably lie in a relatively large window 0.1-100 GeV. However, arguably the most natural mass region is 1-10 GeV.  Indirect detection signals of DM can arise if the DM decays or annihilates producing cosmic rays containing high energy photons, electrons, positrons or antibaryons. Usually, the event density associated to signals of decaying DM depends linearly on the DM density $n_{\\rm DM}$, whereas for annihilating DM the signal has an $n_{\\rm DM}^2$ dependence. If astrophysical gamma-ray signals from DM were detected, then their profile on the sky would let us determine which process produced them, e.g.~\\cite{Boehm:2010qt,Buchmuller:2012rc} (the propagation of charged cosmic rays is affected by galactic magnetic fields, so a signal in these channels would not be so helpful). Whilst decay signals of ADM can readily arise, see e.g.~\\cite{Zhao:2014nsa}, in general, we do not expect late-time annihilation signals from ADM, since the symmetric component is assumed to have annihilated early on. However, if the symmetric component is later `regenerated', this may give rise to annihilation signals. Previous proposals along these lines include slow DM-antiDM oscillations~\\cite{Oscillating}, and intermediate-time decays (in the higher DM-mass regime, \\cite{Falkowski:2011xh}). Alternatively, there are scenarios in which an asymmetry plays a role in determining the DM relic density, but the DM itself is not asymmetric~\\cite{alt}. In this paper, we investigate the prospect of regenerating a symmetric component via decays in the low DM-mass regime.  A challenge in building viable models of ADM is obtaining a high enough annihilation cross section so that the relic density is set by the asymmetry and not by a frozen-out symmetric DM component. For this to occur the cosmologically stable states in the dark sector typically must have annihilation cross sections some factor larger than the thermal value ($\\sim 3 \\times 10^{-26} {\\rm cm}^3 {\\rm s}^{-1}$). Experimental limits on this scenario arise from a variety of sources. Direct searches rule out large regions of parameter space where the interaction with the visible sector is through heavy portal states that can be parameterised as effective operators, see e.g.~\\cite{MarchRussell:2012hi,Buckley:2011kk,LEPdm,ColliderEW},\\footnote{For contact operators with preferential coupling to specific SM states, such as neutrinos~\\cite{Shoemaker:2013tda}, muons, or taus, some of the tension with experimental constraints can be alleviated.} although viable models remain if the mediating states are light. Perturbativity also often limits the cross sections that can be obtained in viable models. Additionally, indirect detection constraints are limiting in some regions of parameter space.  Generally, the strongest astrophysical constraints on the DM annihilation cross section (if the population is entirely symmetric) are close to or below the lower cross-section bound from the ADM relic annihilation constraint. In particular, if a sizeable symmetric component is regenerated before the recombination era, then the limits on the annihilation cross section from CMB observations \\cite{CMB} are typically only a small factor above the thermal freeze-out value, for DM masses $\\lesssim 20~{\\rm GeV}$. Since the CMB limits are mostly a function of energy injection, rather than the particular annihilation products, and are not subject to uncertainties regarding DM distribution, they leave very little parameter space for early decays, for annihilation cross-sections large enough to meet the ADM relic annihilation requirements.   One manner of circumventing these constraints is for the annihilation cross section at late times to be suppressed relative to earlier times, e.g.\\ via velocity-suppression. If the symmetric population is regenerated after freeze-out, but early relative to astrophysical timescales, then the phenomenology of this scenario is much like symmetric DM with a velocity-suppressed cross section, although with the annihilation rate decoupled from the thermal value. An alternative possibility, with distinctive phenomenology, is the case in which the DM remains predominantly asymmetric up to the present time (or at least through the recombination era), with only a small symmetric component regenerated. If the decay timescale is significantly longer than the time at recombination, then CMB constraints are significantly weakened, which can open the possibility of present-day annihilation signals without corresponding CMB signals. Another interesting feature of regenerating the symmetric component through late decays, which we will discuss in some detail, is that its density profile may be modified significantly by the velocity kick from the decay process. Obtaining a sufficiently small velocity kick to keep the regenerated population concentrated in galaxies motivates DM models with a small mass splitting between different components. For these models, the annihilation profile may be different to those arising from conventional models, and could be entirely absent from systems with low escape velocities such as dwarf spheroidal galaxies.  As mentioned, coupling ADM to the Standard Model (SM) without violating direct search constraints requires some model-building, and we consider two scenarios in detail. In one case the annihilation portal to the SM is through a pseudoscalar; this alleviates the direct detection constraints, as the DM-proton scattering cross section is significantly suppressed. If the pseudoscalar mass is just above the value needed for an s-channel annihilation resonance, then annihilations during the freeze-out process (when temperatures are high) are enhanced relative to annihilations at later times. This is essentially an example of the velocity dependence discussed above. The second scenario we consider is where annihilation occurs to two on-shell hidden sector vector bosons, which decay relatively slowly to SM states. The small coupling of the vectors to the SM suppresses direct search constraints, and the rate at which we regenerate the symmetric component controls the strength of CMB and late-time annihilation constraints. These models act as a {\\em proof of principle} that annihilation signals can be generated from certain models of ADM, but the constraints that we discuss are more generally relevant to other models of annihilating DM in the low mass region. The observation of signals in this mass region may suggest that the DM relic density is set by an asymmetry, and that the hidden sector dynamics is correspondingly more complicated than often assumed.  Turning to the structure of the paper, in Sect.~\\ref{S1} we begin by discussing an example class of models that regenerates a symmetric component through decays. Subsequently, in Sect.~\\ref{S2} we explore the constraints on models of annihilating ADM and identify the parameter space for which successful models can be constructed. Additionally, we study the modification of signal morphologies due to the decays. Sect.~\\ref{S3} explores the hidden sector model-building necessary to satisfy the constraints obtained. ",
        "conclusions": "\\label{S4}   ADM is a well motivated framework in which to examine proposals of light DM, as if the baryon and DM asymmetries are comparable, then the fact that $\\Omega_{\\rm{\\rm DM}}\\sim5\\Omega_{B}$ implies that the DM mass is similar to the proton. Annihilation-like indirect detection signatures are typically not expected if the DM relic density is determined by a particle asymmetry, but can arise in models with more complex hidden sectors. We have outlined a new class of ADM models in which annihilation signals are generated, and presented a number of concrete examples. In particular we discussed a scenario involving two states $B',~L'$ with comparable relic densities, stabilised by two different approximately conserved quantum numbers. Subsequently, processes which violate these approximate symmetries lead to decays of the heavier state $B'$, regenerating the symmetric component of the lighter species $L'$.  This allows for the prospect of observable indirect detection signals with annihilation-like profiles via $L'$-pair annihilation.  One of the principal model-building challenges for producing observable signals in this manner is that, unless the mass splitting between $B'$ and $L'$ is small, most of the $\\overline{L'}$ generated via $B'$ decays are immediately ejected from the galaxy and do not give rise to annihilation-type signals. The desire for observable annihilation signals constrains the parameter space of the model as discussed in Sect.~\\ref{sec:splitting} and illustrated in Fig.~\\ref{fig:nfwAlpha}. Moreover, if the mass splitting is moderate then the symmetric component of $L'$ may escape galactic structures with low escape velocities. Thus a distinctive signature of this class of models is that annihilation signals could be observed in our galaxy and conspicuous by their absence in dwarf galaxies. Further, we outlined two scenarios in which such small mass splittings can arise, in the context of broken flavour symmetries, and composite models analogous to the proton-neutron mass splitting.  There is a range of further experimental constraints on these models both from direct and indirect probes, as encapsulated in Figs.~\\ref{fig:regions}, \\ref{fig:decSM}, \\ref{fig:shell10}. In order to evade the strong direct detection bounds we studied two particular scenarios: in the first case annihilations occur on resonance and in the alternative the DM annihilates to pairs of meta-stable hidden sector states. We highlighted the prospect for indirect detection signals in the near future, see for instance Figs.~\\ref{fig:regions} \\& \\ref{fig:pbar}, and commented on tentative signals of DM annihilations near the galactic centre \\cite{Hooper:2013rwa,GC} in Sect.~\\ref{3.6}.  To conclude, we find that whilst models of low-mass ADM with a symmetric component regenerated by decays can produce observable annihilation signals, and satisfy the various experimental bounds, these models are typically required to possess some specific properties and thus exhibit some predictive features. Several experiments have hinted at the possibility of DM in the 1-50 GeV range, both direct detection \\cite{DD} and indirect signals of DM annihilation \\cite{Hooper:2013rwa,GC}. The confirmation of annihilation signals consistent with light DM (possibly with an annihilation cross section different from the thermal freeze-out value) could be an indication of the class of annihilating ADM models proposed here."
    },
    "1402/1402.6727_arXiv.txt": {
        "abstract": "{Methanol masers at 6.7~GHz are associated with high-mass star-forming regions (HMSFRs) and often have mid-infrared (MIR) counterparts characterized by extended emission at 4.5$\\mu$m, which likely traces outflows from massive young stellar objects (MYSOs).} {Our objectives are to determine the milliarcsecond (mas) morphology of the maser emission and to examine if it comes from one or several candidate MIR counterparts in the clusters of MYSOs.} {The European VLBI Network (EVN) was used to image the 6.7\\,GHz maser line with $\\sim$2\\farcm1 field of view toward 14 maser sites from the Torun catalog. Quasi-simultaneous observations were carried out with the Torun 32\\,m telescope.} {We obtained maps with mas angular resolution that showed diversity of methanol emission morphology: a linear distribution (e.g., G37.753$-$00.189), a ring-like (G40.425$+$00.700), and a complex one (e.g., G45.467$+$00.053). The maser emission is usually associated with the strongest MIR counterpart in the clusters; no maser emission was detected from other MIR sources in the fields of view of 2\\farcm1 in diameter. The maser source luminosity seems to correlate with the total luminosity of the central MYSO. Although the Very Long Baseline Interferometry (VLBI) technique resolves a significant part of the maser emission, the morphology is still well determined. This indicates that the majority of maser components have compact cores.} {} ",
        "introduction": "The 6.7\\,GHz methanol maser emission is widely assumed to be associated with massive young stellar objects (MYSOs), but it is still unclear which structures it probes in the circumstellar environment. High angular resolution studies have revealed quite diverse morphologies of methanol maser sources {from simple and linear to curved and complex, or even circularly symmetric ones (e.g., Minier et al.~\\cite{minier00}; Dodson et al.~\\cite{dodson04}; Bartkiewicz et al.~\\cite{bartkiewicz09}; Pandian et al.~\\cite{pandian11}; Fujisawa et al.~\\cite{fujisawa14}). It has been suggested that linear maser structures with velocity gradients indicate circumstellar disks seen edge-on (Norris et al.\\,\\cite{norris98}; Minier et al.~\\cite{minier00}), and the velocity gradients within individual maser clouds perpendicular to the major axis of maser distribution point planar shocks (Dodson et al.\\,\\cite{dodson04}). Arched and ring-like structures can be explained by models of rotating and expanding disks or outflows where the maser arises at the interface between disk/torus and a flow (Bartkiewicz et al.~\\cite{bartkiewicz09}; Torstensson et al.~\\cite{torstensson11}). Detailed proper motion studies of methanol emission, done in only a few sources so far, revealed different scenarios of ongoing phenomena. Observations of the 3D velocity field of methanol masers in the protostellar AFLG~5142 proved that the emission arises in the infalling gas of a molecular envelope with a 300\\,AU radius (Goddi et al.\\,\\cite{goddi11}). Sanna et al.~(\\cite{s10a}) noticed rotation of methanol masers around a central mass in G16.59$-$0.05. Towards G23.01$-$0.41 a composition of slow radial expansion and rotation motions were detected (Sanna et al.~\\cite{s10b}). In IRAS~20126$+$4104 methanol maser spots are associated with the circumstellar disk around the object, and also trace the disk at the interface with the bipolar jet (Moscadelli et al.~\\cite{m11a}).  Cyganowski et al.~(\\cite{cyganowski09}) found that 6.7~GHz methanol masers frequently appear toward mid-infrared (MIR) sources with extended emission at 4.5$\\mu$m, the so-called {\\it extended green objects} (EGOs, Cyganowski et al.\\,\\cite{cyganowski08}) or {\\it green fuzzies} (Chambers et al.~\\cite{chambers08}). This emission probably signposts shocked molecular gas, mainly H$_2$ and CO molecules, related to protostellar outflows. However, detailed studies of each target may also identify its \"falsely appearing {\\it green}\" (De Buizer \\& Vacca~\\cite{debuizer10})}. Simple analysis of the {\\it Spitzer} GLIMPSE data (Benjamin et al.\\,\\cite{benjamin03}; Fazio et al.\\,\\cite{fazio04}) maps proved that the majority of methanol masers from Bartkiewicz et al.~(\\cite{bartkiewicz09}) are spatially associated within 1\\arcsec\\, with EGOs and there is a trend that the regular maser structures coincide with the strongest MIR objects.  In this paper we report on the European VLBI Network (EVN) observations of a sample of 6.7~GHz maser sources. Our aims are threefold: i) to obtain the accurate absolute position of the sources and to determine their morphology, ii) to check if the emission is associated with a single or with several MYSOs, and iii) to determine if the maser morphology is consistent with an outflow. With the new data we discuss the range of physical parameters of central objects which power the methanol maser emission. ",
        "conclusions": "We successfully imaged the 6.7~GHz methanol emission toward 15 targets (four of them belong to the complex HMSFR W49N). Using a short correlation time (0.25~s) we were able to image significantly offset emission of $-$54'' in RA and $-$56'' in Dec (G43.149$+$00.013) from the pointing center. Although the VLBI resolves some of the emission, there is no problem studying the morphology; almost all maser clouds must have compact cores. The maser emission was always associated with the strongest MIR counterpart in the clusters, so we conclude that the appearance of the maser emission is related to the IR brightness of the central MYSO. The maser (isotropic) luminosity and the total luminosity of the central object are likely to correlate; however, a larger sample is needed to verify this relation. The maser linear extent is related to its spectrum width. We also note that the spectra of methanol rings are the widest among the sample and their emission also appears more extended on the sky."
    },
    "1402/1402.6661_arXiv.txt": {
        "abstract": "In this paper, we investigate a scenario of variable gravity and apply it to the unified description of inflation and late time cosmic acceleration dubbed quintessential inflation. The scalar field called ``cosmon\" which in this model unifies both the concepts reduces to inflaton at early epochs. We calculate the slow-roll parameters, the Hubble parameter at the end of inflation, the reheating temperature,the tensor-to-scalar ration and demonstrate the agreement of the model with observations and the Planck data. As for the post inflationary dynamics, cosmon tracks the background before it exits the scaling regime at late times. The scenario gives rise to correct epoch sequence of standard cosmology, namely, radiative regime, matter phase and dark-energy. We show that the long kinetic regime after inflation gives rise to enhancement of relic gravity wave amplitude resulting into violation of nucleosynthesis constraint at the commencement of radiative regime in case of an inefficient reheating mechanism such as gravitational particle production. Instant preheating is implemented to successfully circumvent the problem. As a generic feature, the scenario gives rise to a \u0091blue\u0092 spectrum for gravity waves on scales smaller than the comoving horizon scale at the commencement of the radiative regime. ",
        "introduction": "Theoretical and observational consistencies demand that the standard model of Universe be complemented by early phase of accelerated expansion, dubbed inflation \\cite{Starobinsky:1980te,Starobinsky:1982ee,Guth:1980zm,Linde:1983gd,Linde:1981mu,Liddle:1999mq,Langlois:2004de,Lyth:1998xn,Guth:2000ka,Lidsey:1995np,Bassett:2005xm,Mazumdar:2010sa,Wang:2013hva,Mazumdar:2013gya}, and late time cosmic acceleration \\cite{Copeland:2006wr,Sahni:1999gb, Frieman:2008sn,Padmanabhan:2002ji,Padmanabhan:2006ag, Sahni:2006pa,Peebles:2002gy,Perivolaropoulos:2006ce, Sami:2009dk,Sami:2009jx,Sami:2013ssa}. Inflation is a remarkable paradigm, a single simple idea which addresses logical consistencies of hot Big Bang and provides   a mechanism for primordial perturbations needed to seed the structures in the Universe. As for late time cosmic acceleration, it is now accepted as an observed phenomenon though its underlying cause still remains to be obscure, whereas similar confirmation of inflation is still awaited. Thus, the hot Big Bang and the two phases of accelerated expansion is a theoretically accepted framework for the description of our Universe.  No doubt that inflation is a great idea, the phenomenon should therefore live for ever such that the late time cosmic acceleration is nothing but its reincarnation {\\it \\`a la} quintessential inflation \\cite{Peebles:1999fz,Sahni:2001qp,Sami:2004xk,Copeland:2000hn,Huey:2001ae, Majumdar:2001mm,Dimopoulos:2000md,Sami:2003my, Dimopoulos:2002hm,Dias:2010rg,BasteroGil:2009eb,Chun:2009yu,Bento:2008yx, Matsuda:2007ax,da Silva:2007vt,Neupane:2007mu,Dimopoulos:2007bp, Gardner:2007ib,Rosenfeld:2006hs,Bueno Sanchez:2006ah,Membiela:2006rj, Bueno Sanchez:2006eq,Cardenas:2006py,Zhai:2005ub,Rosenfeld:2005mt, Giovannini:2003jw,Dimopoulos:2002ug,Nunes:2002wz,Dimopoulos:2001qu, Dimopoulos:2001ix,Yahiro:2001uh,Kaganovich:2000fc,Peloso:1999dm, Baccigalupi:1998mn,Hinterbichler:2013we}. The idea was first proposed by Peebles and Vilenkin in 1999 \\cite{Peebles:1999fz} which was later implemented in the framework of braneworld cosmology \\cite{Sahni:2001qp,Sami:2004xk,Copeland:2000hn,Huey:2001ae, Majumdar:2001mm}. At the theoretical level, it sounds pretty simple to implement such a proposal in the language of a single scalar field. The field potential should be shallow at early times, facilitating slow roll, followed by steep behavior thereafter and turning  shallow again at late times. The steep potential is needed for radiative regime to commence, such that the field is sub-dominant during radiation era and does not interfere with nucleosynthesis. It should continue to remain in hiding during matter phase, till its late phases, in order not to obstruct structure formation. It is then desirable to have scaling regime, in which the field mimics the background being invisible, allowing the dynamics to be free from initial conditions, which in turn require a particular steep behavior of the potential. At late times the field should overtake the background, giving rise to late-time cosmic acceleration, which is the case if slow roll is ensured or if the potential mimics shallow behavior effectively.  There are several obstacles in implementing the above unification scheme. First, since  inflation survives in this scenario until late times, the potential is typically of a run-away type and one therefore requires an alternative mechanism of reheating in this case. One could invoke reheating due to gravitational particle production after inflation \\cite{Kofman:1994rk,Dolgov:1982th,Abbott:1982hn, Ford:1986sy,Spokoiny:1993kt,Kofman:1997yn,Shtanov:1994ce,Campos:2002yk}, which is a universal phenomenon. However, the latter is an inefficient process and it might take very long for radiative regime to commence. Clearly, in this case, the scalar field spends long time in the kinetic regime such that the field energy density redshifts with the scale factor as $a^{-6}$ corresponding to equation of state of stiff matter. It is known that the amplitude of gravitational waves produced at the end of inflation enhances during kinetic regime, and if the latter is long, the relic gravitational waves \\cite{Sahni:2001qp,Sami:2004xk,Grishchuk:1974ny,Grishchuk:1977zz, Starobinsky:1979ty,Sahni:1990tx,Souradeep:1992sm, Giovannini:1998bp,Giovannini:1999bh,Langlois:2000ns,Kobayashi:2003cn, Hiramatsu:2003iz,Easther:2003re,Brustein:1995ah,Gasperini:1992dp, Giovannini:1999qj,Giovannini:1997km,Gasperini:1992pa,Giovannini:2009kg, Giovannini:2008zg,Giovannini:2010yy,Tashiro:2003qp} might come into conflict with nucleosynthesis constraint at the commencement of radiative regime \\cite{Sahni:2001qp,Sami:2004xk,Tashiro:2003qp}. Hence, one should look yet for another alternative reheating mechanism, such as instant preheating \\cite{Felder:1998vq,Felder:1999pv,Campos:2004nc}, to circumvent the said problem.  A second obstacle to the unification is that if we want the scalar field to mimic the background for most of the thermal history, the field potential should behave like a steep exponential potential at least approximately such as the inverse power-law potentials. Since scaling regime is an attractor in such cases, an exit mechanism from scaling regime to late time acceleration should be in place in the scenario.  Let us examine as how to build the unified picture. The single scalar field models aiming for quintessential inflation can be broadly put into two classes: (1) Models in which the field potential has a required steep behavior for most of the history of universe but turn shallow at late times, for instance, the inverse power-law potentials \\cite{Sahni:2001qp,Sami:2004xk,Copeland:2000hn,Sami:2003my}. (2) Models in which the field potential is shallow at early epochs giving rise to inflation, followed by the required steep behavior.  In the first class of potentials, we can not implement inflation in the standard framework, since slow roll needs to be assisted in this case. For example, one could invoke Randall-Sundrum (RS) braneworld \\cite{Randall:1999ee,Randall:1999vf} corrections \\cite{Sahni:2001qp,Sami:2004xk,Copeland:2000hn} to facilitate inflation with steep potential at early epochs. In this case, as the field rolls down to low energy regime, the braneworld corrections disappear, giving rise to a graceful exit from inflation and thereafter the scalar field has the required behavior. However, gravitational particle production \\cite{Ford:1986sy,Spokoiny:1993kt} is extremely inefficient in the braneworld inflation \\cite{Sahni:2001qp} and one could in principle introduce the instant preheating to tackle the relic gravitational waves problem \\cite{Sami:2004xk}. Unfortunately, the steep braneworld inflation is inconsistent with observations, namely, the tensor-to-scalar ratio of perturbations is too high in this case. Thus, the scenario fails in the early phase, although the late-time evolution is compatible with theoretical consistency and observational requirements \\cite{Sahni:2001qp,Sami:2004xk}.  In the second class of potentials, that is shallow at early epochs followed by steep behavior, we need a mechanism to exit from scaling regime. A possible way out is provided by introducing neutrino matter, such that neutrino masses are field-dependent \\cite{Wetterich:2013aca,Wetterich:2013jsa,Wetterich:2013wza}. Such a scenario can be motivated from Brans-Dicke framework, with an additional assumption on the matter Lagrangian in the Jordan frame, namely treating massive neutrinos differently from other forms of matter in a way that the field is minimally coupled to cold dark matter/baryon matter in Einstein frame whereas the neutrino masses grow with the field \\cite{Wetterich:2013jsa}. In such a scenario neutrinos do not show up in radiation era; their energy density tracks radiation being sub-dominant. However, in the subsequent matter phase at late times, as they become non-relativistic, their masses begin to grow and their direct coupling to scalar field builds up such that the effective potential acquires a minimum at late times giving rise to late time acceleration, provided the field rolls slowly around the effective minimum. At this point, a question arises, namely whether we could do without neutrino matter and the extra assumption in which case the field would couple to matter directly in the Einstein frame and the effective potential would also acquire a minimum. For simplicity let us assume that we are dealing with a constant coupling $Q$ {\\it \\`a la} coupled quintessence \\cite{Amendola:1999er}. In that case it is possible to achieve slow roll around the minimum of the effective potential, provided that $Q$ is much larger than the slope of the potential, such that the effective equation-of-state parameter has a desired negative value ($w_{\\rm eff}=(-Q+\\al)/(Q+\\al)$ where $\\al$ is the slope of the potential). The scaling solution (which is accelerating thanks to non-minimal coupling), an attractor of the dynamics, is approached soon after the Universe enters into matter-dominated regime and consequently we cannot have a viable matter phase in this case. It is therefore necessary that the matter regime is left intact and the transition to accelerated expansion takes place only at late times. The latter can be triggered by massive neutrino matter with field-dependent masses \\cite{Amendola:2007yx,Wetterich:2007kr,Pettorino:2010bv}.  In this paper we consider a scenario of quintessential inflation in the framework of variable gravity model \\cite{Wetterich:2013jsa,Wetterich:2013aca,Wetterich:2013wza, Wetterich:2014eaa,Wetterich:2014bma}. We first revisit the model in Jordan frame (Sec.~\\ref{JF}) and then we transit to the Einstein frame (Sec.~\\ref{EF}) for detailed investigations of cosmological dynamics by considering the canonical form of the action (Sub Sec.~\\ref{Canonical}). Behavior of the canonical field with respect to the non canonical field is also examined (Sub Sec.~\\ref{sig_behavior}). In the Einstein frame we examine the inflationary phase (Sec.~\\ref{Inflation}), kinetic regime and late time transition to dark energy (Sec.~\\ref{Late times evolution}). Ref.\\cite{Wetterich:2013jsa} provides broad out line of inflation and late time acceleration in the framework of model under consideration. In this paper, we present complete evolution history by invoking suitable preheating mechanism. We investigate issues related to the spectrum of relic gravity waves (Sub Sec.~\\ref{RGW}) as a generic observational features of quintessential inflation. The relic gravity wave amplitude is defined by the inflationary Hubble parameter whereas the spectrum of the wave crucially depends upon the post inflationary evolution. We investigate the problems related to the long kinetic regime in the scenario and discuss the instant preheating (Sub Sec.~\\ref{Instant Preheating}) to tackle the problem. Post inflationary evolution (Sub Sec.~\\ref{Cosmological equations}) is investigated with canonical action and the epoch sequences (Sub Sec.~\\ref{Transient Behavior}) are achieved with viable matter phase. Detailed dynamical analysis is performed to check the nature of stability of all fixed points (Sub Sec.~\\ref{Asymptotic behavior}). Finally in Sec.~\\ref{Conclusions} we summarize the results. ",
        "conclusions": "\\label{Conclusions}  In the present work we have investigated a scenario of variable gravity \\cite{Wetterich:2013jsa,Wetterich:2013aca} in context with quintessential inflation$-$ a unified description of cosmic  evolution from inflation to radiation, matter and dark-energy epochs. In variable gravity the Planck mass is driven by a scalar-field, which additionally drives the mass of the various particles. This field-depending mass, amongst others, leads to the appearance of an effective interaction between the scalar field and matter and neutrinos. Furthermore, through suitable conformal transformations, one can formulate this model in the Einstein frame  in terms of a canonical scalar field with an effective non-minimal coupling between the canonical field and the neutrinos. The cold dark matter is minimally coupled in this framework. The key assumption in the model is related to the field dependence of masses in Jordan frame such that cold dark matter and baryonic matter has standard behavior in Einstein frame whereas the neutrino masses grow with field in a specific way. The canonical scalar field at early times is shown to drive inflation with required number of e-folds $\\mathcal{N}$ (which is approximately equal to $70$ in the model under consideration) and tensor to scalar ratio of perturbations, $r\\simeq 0.11$ consistent with Planck data within $2\\sigma$ confidence level. After inflation, the field potential fast turns into a steep exponential potential such that the field enters into kinetic regime with field energy density redshifting as $a^{-6}$. We checked that gravitational particle production as a reheating mechanism is inefficient($(\\rho_\\phi/\\rho_{r})_{\\rm end}\\simeq 10^{11}$) and it takes long for the radiative regime to commence. The amplitude of relic gravity waves enhances during kinetic regime such that nucleosynthesis constraint($(\\rho_\\phi/\\rho_{r})_{\\rm end}\\lesssim 10^9$) is violated at the beginning of radiation domination in this case. We then implemented instant preheating mechanism which involves the coupling of inflaton with a scalar field $\\chi$ such that $m_\\chi=g|\\phi|$ which in turn couples to matter field, $h\\chi \\bar{\\psi}\\psi$. At the end of inflation, the mass of $\\chi$ changes non-adiabatically giving rise to $\\chi$ production. Assuming that the energy of $\\chi$ production is thermalized,  we can achieve $(\\rho_\\phi/\\rho_{r})_{\\rm end}\\lesssim 10^9$ provided that $g\\gtrsim 6\\al\\times 10^{-5}$. Since after inflation, $\\phi$ grows fast, the produced particles are shown to decay fast into $\\bar{\\psi}\\psi$ avoiding any back reaction of $\\chi$ particles on the post inflationary dynamics of field $\\phi$ provided that $h\\gtrsim 2g^{-1/2}10^{-6}$. We also noticed that particle production takes place almost instantaneously after the end of inflation($\\phi\\leq \\phi_p\\simeq 4\\times 10^{-4}$). We have shown that instant preheating takes place in a large parameter space (g,h) and the process is quite efficient to comply with the thermal history.  Since the field potential in post inflationary era mimics a steep exponential potential with chosen slope, the field exhibits the scaling behavior after the locking regime being sub-dominant. At late times when neutrinos become non-relativistic, the direct coupling of neutrino matter with scalar field builds up and thanks to non-minimal coupling, the field potential acquires minimum which slowly evolves with expansion of universe. The field settles in that minimum for ever; the transition from scaling regime to late time cosmic acceleration is successfully triggered by growing neutrino matter. By performing a detailed phase-space analysis, we showed that apart from the usual stable attractors similar to those of standard quintessence, namely the de Sitter, the dark-energy dominated quintessence-like, and the non-accelerating universe, stiff-dark ones, the system can result in a new attractor, with properties depending on the neutrino behavior, corresponding to a quintessence-like universe.  We have shown that quintessential inflation based upon variable gravity model can successfully unify inflation and dark energy. The model based upon instant preheating is shown to be consistent with observations. The possibility of  detection of the relic gravity wave background by Advanced LIGO  and LISA are discussed.  As for early universe, the model complies with the recent Planck data, showing that within $2\\sigma$  confidence level, the scenario is consistent with observations. The scrutiny of late time acceleration and the study of observational constraints on the model parameters is deferred to our future work. It will also be interesting to carry out detailed investigations of stability of neutrino matter under perturbations. One may also examine the scenario under consideration in the framework of warm inflation which might further improve the tensor to scalar ratio of perturbations."
    },
    "1402/1402.5986_arXiv.txt": {
        "abstract": "We examine a possible formation scenario of galactic thick discs with numerical simulations. Thick discs have previously been argued to form in clumpy disc phase in the high-redshift Universe, which host giant clumps of $\\hspace{0.3em}\\raisebox{0.4ex}{$<$}\\hspace{-0.75em}\\raisebox{-.7ex}{$\\sim$}\\hspace{0.3em}10^{9}~{\\rm M_\\odot}$ in their highly gas-rich discs. We performed $N$-body/smoothed particle hydrodynamics simulations using isolated galaxy models for the purpose of verifying whether dynamical and chemical properties of the thick discs formed in such clumpy galaxies are compatible with observations. The results of our simulations seem nearly consistent with observations in dynamical properties such as radial and vertical density profiles, significant rotation velocity lag with height and distributions of orbital eccentricities. In addition, the thick discs in our simulations indicate nearly exponential dependence of $\\sigma_\\theta$ and $\\sigma_z$ with radius, nearly isothermal kinematics in vertical direction and negligible metallicity gradients in radial and vertical directions. However, our simulations cannot reproduce altitudinal dependence of eccentricities, metallicity relations with eccentricities or rotation velocities, which shows striking discrepancy from recent observations of the Galactic thick disc. From this result, we infer that the clumpy disc scenario for thick-disc formation would not be suitable at least for the Milky Way. Our study, however, cannot reject this scenario for external galaxies if not all galaxies form their thick discs by the same process. In addition, we found that a significant fraction of thick-disc stars forms in giant clumps. ",
        "introduction": "\\label{Intro} Duplex structures of thin and thick discs seem to be ubiquitous among disc galaxies \\citep[e.g.][]{b:79,db:02,yd:06}, including the Milky Way (MW) \\citep[e.g.][]{y:82,gr:83}. Because thick-disc stars are generally old and $\\alpha$-enhanced population which implies their early and rapid formation \\citep[e.g.][]{abw:06,hdl:13}, structures of the thick discs are expected to reflect galactic formation history. The origin of the thick discs is, however, still an open question, and a number of possible formation scenarios have been proposed so far. Ones frequently discussed in recent studies are as follows: stellar accretion from disrupted satellites \\citep[e.g.][]{ans:03}, thickening of a thin disc by minor mergers \\citep[e.g.][]{qhf:93,kbz:08}, multiple gas-rich minor mergers \\citep{bkg:04,bgm:05,bsg:12}, and radial migration by spiral arms and barred structures in a disc (e.g. \\citealt{sb:09b,sb:09a}). \\citet{sha:09} have discussed that distribution of orbital eccentricities of thick-disc stars reflects their formation processes and can be used as a probe of these scenarios. Afterwards, some observations were performed for the thick disc of the MW and elucidated that the thick-disc stars have a mono-modal distribution with the peak at the eccentricity of $\\simeq0.2$ \\citep{dkr:10,lba:11,dgk:11,r:11,krl:11,kgw:13}. They argued that the multiple gas-rich minor mergers would be a favourable scenario for the Galactic thick-disc formation, and the radial migration scenario could also reproduce the eccentricity distribution in acceptable matching. On the other hand, the stellar accretion scenario results in relatively higher eccentricities than the observations, and the thin-disc thickening by dwarf galaxies seems to bring about a bi-modal distribution. In addition, observations of \\citet{lba:11} found that correlations between metallicities and some kinematic properties are obviously different between the Galactic thin and thick discs.  In the high-redshift Universe, some galaxies have been observed to be clumpy\\footnote{These discy galaxies are often called clump clusters, tadpole and chain galaxies according to their inclinations.} , and they have giant clumps the masses of which are $\\hspace{0.3em}\\raisebox{0.4ex}{$<$}\\hspace{-0.75em}\\raisebox{-.7ex}{$\\sim$}\\hspace{0.3em}10^{9}~{\\rm M_\\odot}$ in their highly gas-rich discs \\citep[e.g.][]{bar:96,eeh:04,fgl:06,g:06,ggf:11,bce:13,tkt:13II}. Numerical simulations demonstrated the evolution from clump clusters to disk galaxies \\citep[e.g.][]{n:98,n:99,isg:04,bee:07,atm:09,dsc:09,cdb:10,mdc:13}. In the evolutionary process, giant clumps fall into the galactic centre by dynamical friction and can eventually form a galactic bulge \\citep{n:98,n:99,ebe:08,is:12,dk:13,bpr:13,pvt:13}. In addition, the clumps are so massive that they can change dark matter (DM) density profiles (\\citealt{ebe:08b,is:11}; see also \\citealt{dph:14}) and induce significant net rotation of halo stars and globular clusters around galactic discs \\citep{i:13}.  \\citet{n:96} has pointed out the possibility of thick-disc formation in such a clumpy phase of galaxy evolution: Orbital motions of giant clumps kinematically can heat up a galactic disc, and disrupted clumps supply stars to the thick disc to some extent. Afterwards, \\citet{bem:09} demonstrated it by numerical simulations. Their numerical simulations showed that the clumpy disc scenario can reproduce observed thick discs with exponential density profiles \\citep[see also][]{bee:07,es:13} and scale heights independent from radius. Thick discs of external galaxies were indeed observed to have a constant thickness \\citep{vs:81,jt:82,sg:90,matthews:00,db:02}. They suggested that the reproducibility of the constant thickness is a characteristic point of this formation scenario; the thin-disc thickening by minor mergers lead to flaring of thick discs in outer disc regions \\citep[e.g.][]{kbz:08}. Observations also seemed to show that the disc-like structures in the clumpy galaxies are consistent with thick discs in nearby galaxies in their thickness and age \\citep{ee:05,ee:06}, and spatial distributions of clumps are nearly exponential although surface brightness profiles of the discs are clearly not \\citep{eev:05}. Thus, it could be said that the thick-disc formation in clumpy galaxies is also a possible scenario and should be examined in detail. However, the properties of the thick discs formed in this scenario have not been compared with observations in the way of \\citet{sha:09}.  In this study, we performed several runs of high-resolution $N$-body/smoothed particle hydrodynamics (SPH) simulation using isolated models and study dynamical and chemical properties of thick discs formed in the clumpy disc scenario. Our aim in this paper is to scrutinize the credibility of this thick-disc formation scenario and discuss whether we can verify or discard the scenario. We describe our simulation settings in \\S\\ref{sim} and analyses of the thick-disc properties in \\S\\ref{Results}. We present discussion and conclusions in \\S\\ref{discuss} and \\S\\ref{conc}. ",
        "conclusions": "\\label{conc} In this study, we found that properties of thick discs are almost consistent among our runs with different initial conditions. Our simulations indicated that kinematic properties such as density and velocity profiles and distributions of orbital eccentricities are in acceptable agreement with Galactic and external galactic observations. However, altitudinal dependence of the eccentricities and metallicity relations with rotation velocities and eccentricities are not confirmed in our simulations, which are discrepant from recent observations of the Galactic thick disc. Therefore, we infer that the thick-disc formation scenario by clumpy discs would not be suitable at least for the MW."
    },
    "1402/1402.0930_arXiv.txt": {
        "abstract": "We consider a class of leptogenesis models in which the lepton asymmetry arises from dark matter annihilation processes which violate $CP$ and lepton number.  Importantly, a necessary one-loop contribution to the annihilation matrix element arises from absorptive final state interactions. We elucidate the relationship between this one-loop contribution and the $CP$-violating phase. As we show, the branching fraction for dark matter annihilation to leptons may be small in these models, while still generating the necessary asymmetry. ",
        "introduction": "Both the baryon asymmetry of the universe (BAU) and the existence of non-baryonic dark matter (DM) are well motivated by observation at a variety of scales and epochs. Recent observations of the cosmic microwave background (CMB) by the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{WMAP:2011} establish the baryonic and cold dark matter densities in our universe as \\bea \\Omega_b h^2 &\\sim & 0.022  , \\nonumber\\\\ \\Omega_{DM} h^2 &\\sim & 0.12 . \\eea The most common theoretical paradigm used to explain the factor of $ \\sim 5 $ between the two observed densities has been asymmetric dark matter (ADM) models~\\cite{Zurek:2013}. While ADM models transfer an asymmetry of the dark sector into the baryonic sector, dynamical generation of the baryon asymmetry from dark matter decays have also been investigated~\\cite{Kohri:2009}. More recently, weakly interacting massive particle (WIMP) annihilation has been suggested as a mechanism for baryogenesis~\\cite{Cui:2012}. In such WIMPy models, a departure from thermal equilibrium is guaranteed by the WIMP framework. As such, these models of DM annihilation need only violations of baryon number, $C$ and $CP$ in order to satisfy the Sakharov conditions for baryogenesis~\\cite{Sakharov:1967dj}.  \\subsection{Matter Asymmetry}  In order to demonstrate the mechanism of $CP$ violation when dynamically generating an asymmetry, consider a generic  DM annihilation to a multi-particle final state, $XX \\rightarrow Y$, and the $CP$ conjugate process $XX \\rightarrow \\bar Y$. We can write the matrix element for such processes as the sum of $CP$-invariant and $CP$-violating terms,  \\bea {\\cal M}_{XX \\rightarrow Y} &=& {\\cal M}_{XX \\rightarrow Y}^{CP} + {\\cal M}_{XX \\rightarrow Y}^{CPV} , \\nonumber\\\\ {\\cal M}_{XX \\rightarrow \\bar Y} &=&\\pm \\left( {\\cal M}_{XX \\rightarrow Y}^{CP} - {\\cal M}_{XX \\rightarrow Y}^{CPV} \\right) . \\eea  As a result of the optical theorem, the relative phase between the $CP$-invariant and $CP$-violating terms will be $\\pm \\pi /2$ at tree level, eliminating any asymmetry. Thus, any nonvanishing asymmetry in the production cross section must be an interference of $CP$-invariant and $CP$-violating matrix elements with a relative phase difference from an additional imaginary matrix element contribution at loop level,  \\bea \\sigma_{XX \\rightarrow Y} - \\sigma_{XX \\rightarrow \\bar Y} &\\propto & Re \\left[{\\cal M}_{XX \\rightarrow Y}^{CP} ({\\cal M}_{XX \\rightarrow \\bar Y}^{CPV} )^* \\right]. \\eea  In standard models of baryogenesis the relative phase is introduced by the interference of a $CP$-invariant tree-level diagram with a $CP$-violating loop-level diagram. Alternatively, we separate the $CP$-violating phase of two tree-level diagrams from an absorptive phase introduced by a final state decay.  \\subsection{WIMPy Baryogenesis Models}  Before discussing a particular WIMPy model, some general features and issues are worth noting. A pair of WIMPs, stabilized by a discrete symmetry, will annihilate to a quark (or lepton) and a heavy exotic field through a set of operators with the $C$ and $CP$ properties necessary to satisfy the Sakharov conditions. This new state needs to also be protected from interactions with standard model by a discrete charge and should transfer its own asymmetry into a hidden sector. In models with a relative phase generated by the interference of tree-level and loop-level diagrams, dangerous washout diagrams arise. Washout processes yield terms in the Boltzmann equations for the rate of the asymmetry production that are proportional to the asymmetry itself. As a consequence, these $CP$-invariant terms will dampen the generation of an asymmetry.   \\begin{figure}[htb] \\centering \\includegraphics[width=\\textwidth]{loops2-eps-converted-to.pdf} \\caption{Asymmetry generating DM, $X$, annihilation diagrams and corresponding washout diagrams in WIMPy models~\\cite{Cui:2012}. Note that in order to generate a non-vanishing asymmetry, DM must annihilate to a SM lepton (or quark), $L$, and an additional weak scale field, $\\psi$, through the interference of tree-level (top left) and loop-level (top right) diagrams. The loop diagram implies the existence of the leading order, ``pure\", washout process (bottom left) and crossing the tree-level DM annihilation implies the ``mixed\" washout diagram (bottom right). } \\label{fig:loops} \\end{figure}  The presence of these dangerous washout processes are clearly implied by the loop-level DM annihilation diagram in Figure~\\ref{fig:loops}. Because there is no dark matter in the initial state for these ``pure\" washout processes, they lack the Boltzmann suppression present in ``mixed\" washout processes. Thus, the $CP$-violating processes must account for a large portion of the total DM annihilation in order to generate a large enough asymmetry. As the mediator, $S_{1,2}$, runs in the loop, the cross section of the asymmetry generating process will also be inversely proportional to the square of the mediator mass. The result is a necessarily small separation of scales between the new external states, $X$ and $\\psi$, and the mediator $S_{1,2}$. This is generally problematic for models constructed within the framework of an effective field theory (EFT) and for WIMPy models which attempt to independently resolve the DM relic density, BAU and ultraviolet (UV) physics. ",
        "conclusions": "We have shown that the $CP$-violating process required by the Sakharov conditions for the dynamical generation of the BAU can arise from tree-level diagrams and an absorptive final state interaction. This mechanism sequesters the one-loop suppression of the asymmetry and eliminates dangerous tree-level washout processes which are not Boltzmann suppressed. As a result, we can  effectively decouple our dark sector from the UV and address the DM relic density and BAU  more independently. We also have carefully treated nonperturbative sphaleron effects through the electroweak crossover.  For simplicity, we have focused only on the approximation where the intermediate $H$ is a narrow resonance and can be produced on-shell. Although one would expect an even more efficient asymmetry generation if we disregarded this assumption, additional washout processes incapable of producing a narrow resonance would have to be considered. Also, baryogenesis through DM annihilation directly into quarks would avoid the sphaleron shutoff, but the mass of $H$ would be highly constrained by collider searches for DM production by quarks. In principal, the application of final state absorptive interactions could also be used to eliminate dangerous washout processes from more traditional baryogenesis or leptogenesis scenarios."
    },
    "1402/1402.2104_arXiv.txt": {
        "abstract": "{We present a coherent and detailed Monte Carlo simulation of the  population   of  hot  white  dwarfs.   We  assess  the statistical significance of the hot  end of the white dwarf luminosity function  and the role played  by the bolometric corrections of hydrogen-rich white dwarfs at high effective temperatures.} {We use  the most up-to-date stellar  evolutionary models and implement a full description of the observational selection biases  to obtain  realistic  simulations  of the  observed white dwarf population.} {Our  theoretical results  are compared  with the  luminosity function  of  hot  white  dwarfs obtained  from  the  Sloan Digital Sky  Survey (SDSS),  for both  DA and  non-DA white dwarfs.} {We  find  that  the  theoretical  results  are  in  excellent agreement with the observational  data for the population of white  dwarfs with  hydrogen  deficient atmospheres  (non-DA white  dwarfs).  For the  population  of  white dwarfs  with hydrogen-rich atmospheres  (white dwarfs  of the  DA class), our   simulations   show   some   discrepancies   with   the observations for the brightest luminosity bins, namely those corresponding  to   luminosities  larger  than   $\\sim  10\\, L_{\\sun}$. These discrepancies can  be attributed to the way in which the masses of the white dwarfs contributing to this luminosity  bin have  been computed,  as most  of them  have masses  smaller   than  the  theoretical  lower   limit  for carbon-oxygen white dwarfs.} {We  conclude  that  the   way  in  which  the  observational luminosity function of hot white dwarfs is obtained is very sensitive to  the particular  implementation of  the method used to derive the masses of  the sample. We also provide a revised  luminosity  function  for hot  white  dwarfs  with hydrogen-rich atmospheres.} ",
        "introduction": "White dwarfs are  the most common end-point of  stellar evolution.  In fact, more than  $\\sim 95\\%$ of the stars in  our Galaxy, namely those with masses smaller than $\\sim 10\\, M_{\\sun}$, will end their lives as white dwarfs. Hence, the white  dwarf population carries very valuable information about the  evolution of the vast majority of  stars in our Galaxy.  Moreover,  their structural  and evolutionary  properties are relatively well understood -- see,  for instance, the recent review of \\citet{Althausetal10}.   Consequently,  the characteristics  of  white dwarf populations can be used  to constrain the evolutionary models of their progenitor stars.  For instance,  the white dwarf population can be used  to assess  the amount  of mass lost  by their  progenitors, a process that  plays a  key role  in stellar  evolution.  Additionally, since the  mass ejected  in the  late stages  of stellar  evolution is enriched  in heavy  elements, the  study of  the Galactic  white dwarf population is important to model the chemical evolution of our Galaxy. These, however, are  not the only useful applications  of studying the population of white dwarfs.  It is  worth mentioning that the study of the local  white dwarf population provides  interesting constraints on the local star  formation rate \\citep{DPetal94,Rowell} and  on the age of    the     Galactic    disk     in    the     Solar    neighborhood \\citep{winget,nature,GB88} and  other Galactic  populations, including open  \\citep{nature2}  and  globular  clusters --  see,  for  example, \\cite{hansen07}.   Furthermore, the  ensemble  characteristics of  the population of white dwarfs can be used to constrain the mass of weakly interacting  particles, like  axions \\citep{Isern},  neutrino emission rates  \\citep{lambvh75} and  the nature  of the  baryonic dark  matter content of our Galaxy -- see, for instance, \\citet{halo}, \\citet{MC2}, and \\citet{MC3},  and references therein.  Finally,  the population of white dwarfs can also be used to study alternative theories of gravity \\citep{JCAP,gdot}.  All these investigations require  not only accurate evolutionary white dwarf  cooling  sequences, but  also  reliable  determinations of  the observed ensemble properties of the white dwarf population, namely its luminosity function and  its mass distribution.  Thus,  to address all these  questions,  we   need  large  white  dwarf   samples  of  known completeness.   The Sloan  Digital  Sky Survey  (SDSS)  -- see,  e.g., \\cite{yor00} --  has provided  us with  a very  large sample  of white dwarfs and from it, a reliable white dwarf luminosity function for hot white dwarfs -- namely those  with effective temperatures in excess of 23,500~K --  has been  recently obtained \\citep{Krz2009}.   This white dwarf luminosity function has the interesting benefit that it has been obtained  using solely  spectroscopically-confirmed white  dwarfs from the SDSS DR4 and it is thus an excellent testbed to check not only the white  dwarf cooling  sequences  at high  luminosities,  but also  our ability to  model reliably some  Galactic inputs necessary  to compute the  luminosity  function.    Here  we  describe  the   results  of  a comprehensive set  of Monte Carlo  simulations aimed to model  the hot part of  the white  dwarf luminosity  function for  both hydrogen-rich (DA) white dwarfs and non-DA (hydrogen deficient) stars.  \\begin{figure}[t] \\vspace{8.5cm} \\special{psfile=fig01.ps hscale=50 vscale=50 voffset=-100 hoffset=-24} \\caption{White dwarf luminosity functions for  hot DA and non-DA white dwarfs for two  different models -- see  text Sect.~\\ref{sec:LF} for details.   The solid  line  with filled  squares  is the  luminosity function  of \\citet{Krz2009},  while  the dashed  lines with  hollow squares show the results of our simulations.} \\label{f:hotlf} \\end{figure} ",
        "conclusions": "\\begin{table} \\centering \\begin{tabular}{ccc} \\hline \\hline $M_{\\rm bol}$ & $\\log N$~(pc$^{-3}$~$M_{\\rm bol}^{-1}$) & Number of stars \\\\ \\hline \\noalign{\\smallskip} 2    & $-6.937^{+0.230}_{-0.284}$ &  12 \\\\ \\noalign{\\smallskip} 3    & $-6.342^{+0.138}_{-0.179}$ &  45 \\\\ \\noalign{\\smallskip} 4    & $-6.326^{+0.139}_{-0.173}$ &  47 \\\\ \\noalign{\\smallskip} 5    & $-6.018^{+0.076}_{-0.089}$ &  75 \\\\ \\noalign{\\smallskip} 6    & $-5.739^{+0.063}_{-0.072}$ & 115 \\\\ \\noalign{\\smallskip} 7    & $-5.375^{+0.067}_{-0.079}$ & 209 \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\caption{Observed  DA white  dwarf  luminosity  function dwarfs  after applying  the  correction  for   those  objects  with  masses  below $0.49\\,M_{\\sun}$.} \\label{t:wdlf} \\end{table}  We presented  a set of simulations  of the luminosity function  of hot white  dwarfs.   Our  results  are in  excellent  agreement  with  the observations for non-DA stars, while for DA white dwarfs we found that a plateau  present in the  luminosity function at  $M_{\\rm bol}\\approx 2.0$ --  or, equivalently, $\\log  (L/L_{\\sun})\\approx 1$ --  cannot be reproduced by  our simulations. We  first discarded that  this plateau could be due to  a global Galactic input, like a  recent burst of star formation, because if  this were the case, the  luminosity function of hot  non-DA white  dwarfs  should  show a  similar  plateau.  We  then investigated other  possible origins of this  plateau.  In particular, we studied the  possible role of NLTE effects  and metal contamination in  the atmospheres  of hydrogen-rich  white dwarfs,  the role  of the so-called  Balmer-line problem  for  very hot  white  dwarfs, and  the possibility  that  the  available   cooling  sequences  were  severely affected by  the initial conditions  adopted to compute them,  and the statistical uncertainties in the white  dwarf sample. We found that at the  relevant  luminosities  of  the  plateau,  the  effects  of  NLTE corrections are  limited and  that metal  contamination plays  a minor role in shaping  the luminosity function.  This is also  the case when the Balmer-line problem is taken  into account.  We also discarded the explanation that the available cooling tracks are severely affected by the choice of initial conditions.  Thus,  we were left with a possible statistical fluctuation,  which we  investigated in detail.   We found that even though  our theoretical simulations yield  more white dwarfs in  the   luminosity  range   of  the  plateau   than  exist   in  the \\citet{Krz2009} sample,  their contribution  to the number  density of stars in  this region is  smaller than that  of the SDSS  sample, thus pointing towards an artifact in  the derivation of the SDSS luminosity function.  We therefore analyzed the distribution  of masses and distances of the SDSS sample and we found that  within the range of luminosities of the plateau, the SDSS white dwarf luminosity function is almost completely dominated  by the  contributions of  a  handful of  white dwarfs  with masses smaller than $0.49\\, M_{\\sun}$,  the lower theoretical mass for a carbon-oxygen white dwarf. Thus, either these objects are members of unresolved binaries,  or they have erroneously  determined masses.  We found that once these objects are excluded from the calculation of the DA white dwarf  luminosity function, the agreement  with our simulated luminosity  function  is  excellent.  Additionally,  we  computed  the distances setting the mass of each  of these $< 0.49\\, M_{\\sun}$ white dwarfs at  $0.49\\, M_{\\sun}$,  again finding excellent  agreement with our simulated  luminosity functions and  that obtained with  a $0.49\\, M_{\\sun}$ mass cut applied to  the SDSS sample.  We therefore conclude that the luminosity function of  hot DA white dwarfs of \\cite{Krz2009} is contaminated by white dwarfs with incorrectly determined low masses (below  $0.49\\,  M_{\\sun}$)  and  we provide  a  revised  white  dwarf luminosity function in Table~\\ref{t:wdlf}.   The analysis presented in this paper also  confirms the validity of the  white dwarf theoretical cooling sequences at high effective temperatures via the good match of our model  luminosity functions to  that of the SDSS  observations. We also  emphasize that  the way  in which  the observational  luminosity function of  hot white  dwarfs is  obtained is  very sensitive  to the particular implementation of  the method used to derive  the masses of the sample."
    },
    "1402/1402.0101_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "String theory, M-theory, and the related models of brane interactions suggests new and interesting scenarios for the birth of the Universe and its subsequent primordial evolution, not necessarily limited in time by a big bang singularity. They can be tested by present (or near-future) observations concerning the properties -- in particular, the ``blue'' tilt \\cite{29a} of the spectrum --  of the cosmic background of relic gravitational radiation (see, e.g., \\cite{30}), the possibility of axion contributions to the CMB anisotropies (see, e.g., \\cite{30a}), the evolution of the so-called dark energy (or quintessence) field dominating the large-scale dynamics (see, e.g., \\cite{31, 31b}).  \\index{string-gas cosmology} \\index{brane-gas cosmology} Some of those scenario have been briefly introduced and illustrated in the previous sections. But there are also other, equally interesting scenarios closely related to the previous ones, among which I would like to mention the string-gas \\cite{32} and brane-gas \\cite{33} cosmologies, based on the repulsive mechanism of winding modes, as well as more general bouncing cosmology models (see, e.g., \\cite{34,34b,35}). Also, models of brane anti-brane inflation \\cite{36,37}, where the (time-varying) distance between the two branes plays the role of the inflation field. \\index{bouncing cosmologies}  All these models have many (and interesting) phenomenological implications, but -- as usual in a cosmological context -- many studies and many observational data are required before being able of selecting the model most appropriate to our Universe. Thus, we can easily predict that we still have in front of us many years of work and -- maybe -- of surprising findings."
    },
    "1402/1402.2332_arXiv.txt": {
        "abstract": "{About 40\\% of the observation time of the High Energy Stereoscopic System (H.E.S.S.)  is dedicated to studying active galactic nuclei (AGN), with the aim of increasing the sample of known extragalactic very-high-energy (VHE, $E>100{\\rm\\ GeV}$) sources and constraining the physical processes at play in potential emitters.} {H.E.S.S. observations of AGN, spanning a period from April 2004 to December 2011, are investigated to constrain their $\\gamma$-ray fluxes. Only the 47 sources without significant excess detected at the position of the targets are presented.} {Upper limits on VHE fluxes of the targets were computed and a search for variability was performed on the nightly time scale.} {For 41 objects, the flux upper limits we derived are the most constraining reported to date. These constraints at VHE are compared with the flux level expected from extrapolations of {\\it Fermi}-LAT measurements in the two-year catalog of AGN. The H.E.S.S. upper limits are at least a factor of two lower than the extrapolated {\\it Fermi}-LAT fluxes for 11 objects. Taking into account the attenuation by the extragalactic background light reduces the tension for all but two of them, suggesting intrinsic curvature in the high-energy spectra of these two AGN.} {Compilation efforts led by current VHE instruments are of critical importance for target-selection strategies before the advent of the  Cherenkov Telescope Array, CTA.} \\offprints{\\\\Jonathan Biteau - email: biteau(at)in2p3.fr\\\\ David Sanchez - email: david.sanchez(at)lapp.in2p3.fr} ",
        "introduction": "\\label{Intro}   Since the discovery of their extragalactic origin fifty years ago \\citep{1963Natur.197.1040S}, the class of astrophysical sources called active galactic nuclei (AGN) has been a prime target for astronomers that observe the sky from radio wavelengths to very-high-energy $\\gamma$~rays (VHE, $E>100\\ {\\rm GeV}$). AGN are thought to host super-massive black holes (typical mass of $10^{8-9}M_\\sun$) surrounded by an accretion disk, with a fraction of them showing two-sided jets. To unify the various subclasses of AGN, a scheme to sort them based on their orientation with respect to the observer's line of sight has been proposed since the 1990s \\citep{Urry}. Objects whose jets are closely aligned with the line of sight are called blazars. They fall into two source classes, broad-line-spectrum sources called flat spectrum radio quasars (FSRQ), and BL Lac objects (hereafter BL in tables), which show faint lines or featureless spectra.  Active galactic nuclei, in particular blazars, are the most numerous objects detected at high energy (HE, $100\\ {\\rm MeV}<E<100\\ {\\rm GeV}$), where all-sky surveys can be performed with pair-conversion detectors onboard satellites such as the \\textit{Fermi} Large Area Telescope \\citep[LAT,][]{2009ApJ...697.1071A}. The Second LAT AGN Catalog, hereafter 2LAC, comprises 886 off-plane (i.e. above a Galactic latitude of 10\\dgr) point-like sources associated with AGN that were detected in two years of operation \\citep[][]{2LAC}. AGN constitute a third of the sources known at VHE, despite a coverage biased toward Galactic sources. With the fast decrease of fluxes with increasing energy, observations at VHE are mostly performed with ground-based imaging atmospheric Cherenkov telescopes (IACT), which have a field of view (FoV) of a few degrees but an effective area on the order of $10^5\\,{\\rm m}^2$. Their current sensitivity prevents an all-sky scan in a reasonable amount of time, and IACT observations must be pointed to targets of interest (see e.g. \\citealp{2013APh....43..317D} for a discussion of the capabilities of next generation instruments). Targeted AGN are selected based on their radio and X-ray spectra \\citep{1996ApJ...473L..75S,2000AIPC..515...53P,2002AA...384...56C} as well as based on their HE flux extrapolated to VHE \\citep{2010MNRAS.401.1570T}.  The High Energy Stereoscopic System \\citep[\\hess,][]{Crab} has significantly contributed to the expansion of the class of VHE AGN, with the detection of 23 objects, including 20 discoveries, among 56 known sources of this type as of the end of the year 2013\\footnote{TeVCat, http://tevcat.uchicago.edu/}. The \\hess\\ experiment is located in the Khomas Highland, Namibia (23\\dgr16'18'' S, 16\\dgr30'01''E) at an altitude of $\\unit[1800]{m}$ above sea level. In its first phase, this experiment was an array of four identical telescopes with cameras composed of 960 photomultipliers and segmented reflectors paving a reflective area of $\\unit[107]{m^2}$, for an equivalent diameter of $\\unit[12]{m}$. Most of the AGN detected with \\hess\\ belong to the BL Lac class, as shown in Table~\\ref{Table:Detection}, with the exception of the two nearby radio galaxies of Fanaroff-Riley~I type (FR~I) Centaurus~A and M~87, the FSRQ PKS~1510-089, and the blazar candidate HESS J1943+213, which is located in the Galactic plane. In addition to constraining the radiative processes responsible for the $\\gamma$-ray emission (for detailed discussions, see, e.g., the references in Table~\\ref{Table:Detection}), the VHE spectra of these objects can also serve cosmological purposes, as shown with the constraints \\citep{EBLAHA,2007A&A...471..439M} and indirect measurement \\citep{2013A&A...550A...4H} of the extragalactic background light (EBL). With peak intensities in the optical and far-infrared bands, the EBL is composed of the integrated emission of stars and galaxies as well as of the reprocessing of UV-to-optical light by dust. The EBL can hardly be measured directly, although it is the second-most intense diffuse radiation in the Universe after the cosmic microwave background.  \\begin{table*} \\centering \\small \\begin{tabular}{ l c c c } \\hline\\hline Object & $z$ & Type & Reference\\\\ \\hline Cen A & 0.002 & FR I & \\cite{2009ApJ...695L..40A}\\\\ M 87 & 0.004 & FR I & \\cite{2012ApJ...746..151A}\\\\ Markarian 421 & 0.031 & BL & \\cite{MrkHESS2004}\\\\ AP Librae & 0.049 & BL & \\cite{2012IAUS..284..411S} \\\\ PKS 1440-389 & 0.065 & BL & \\cite{2012ATel.4072....1C} \\\\ PKS 0548-322 & 0.069 & BL & \\cite{2010AA...521A..69A} \\\\ PKS 2005-489 & 0.071 & BL & \\cite{2011AA...533A.110H} \\\\ RGB J0152+017 & 0.080 & BL & \\cite{2008AA...481L.103A} \\\\ SHBL J001355.9-185406 & 0.095 & BL & \\cite{SHBL} \\\\ 1ES 1312-423 & 0.105 & BL & \\cite{2013MNRAS.434.1889H} \\\\ PKS 2155-304 & 0.116 & BL & \\cite{2012AA...539A.149H} \\\\ 1ES 0229+200 & 0.140 & BL & \\cite{0229}\\\\ 1RXS J101015.9-311909 & 0.143 & BL & \\cite{2012AA...542A..94H}\\\\ H 2356-309 & 0.165 & BL & \\cite{2010AA...516A..56H}\\\\ 1ES 1101-232 & 0.186 & BL & \\cite{1101}\\\\ 1ES 0347-121 & 0.188 & BL & \\cite{0347}\\\\ PKS 0301-243 & 0.266 & BL & \\cite{2013AA...559A.136H} \\\\ 1ES 0414+009 & 0.287 & BL & \\cite{2012AA...538A.103H}\\\\ PKS 1510-089 & 0.361 & FSRQ & \\cite{1510} \\\\ PKS 0447-439 & $<0.57\\ddagger$ & BL & \\cite{2013AA...552A.118H}\\\\ PG 1553+113 & - & BL & \\cite{2008AA...477..481A}\\\\ HESS J1943+213 & - & - & \\cite{2011AA...529A..49H}\\\\ KUV 00311-1938 & $>0.506\\dagger\\ $ & BL & \\cite{becherini:490} \\\\ \\hline \\end{tabular} \\tablebib{$\\dagger$ see \\citet{2012AIPC.1505..566P}. $\\ddagger$ see \\citet{2013AA...552A.118H}.} \\caption{AGN detected by \\hess\\ as of September 2013. The redshift, classification and latest \\hess\\ publication on the source are given in columns two to four. Acronyms are defined in the text.} \\label{Table:Detection} \\end{table*}   During the eight years of the first phase of \\hess, some of the observations did not result in significant excesses at the position of the target or in the FoV of the telescopes. A first set of upper limits \\citep[][hereafter HUL1]{2005AA...441..465A} on 19 AGN observed during 63 hours was published after two years of observation. A second paper \\citep[][hereafter HUL2]{2008AA...478..387A} listed 14 upper limits based on 94 hours of observation spanning 2005-2007. In this paper, which follows extensive compilation efforts from previous-generation instruments such as Whipple \\citep{2004ApJ...603...51H} or HEGRA \\citep{2004AA...421..529A}, 47 selected candidates are studied, with observations spanning April 2004 to December 2011, for a total live time of approximately 400 hours. The candidates and the data selection are presented in Sec.~\\ref{Obs}, the event and spectral analyses are examined in Sec.~\\ref{Ana}, and the constraints on the VHE emission are discussed in Sec.~\\ref{Disc}, together with the target-selection strategy. ",
        "conclusions": "\\label{Ccl}  A large sample of AGN has been observed with \\hess\\ since 2002, resulting in the discovery of more than a third of the known extragalactic VHE emitters. Observations of 47 targets without significant excess were selected and upper limits on their integral fluxes were computed. For 41 of these objects, the upper limits derived in this paper are the strongest to date.  No significant flaring event was detected during the $\\sim\\unit[400]{h}$ of observation of the 47 targets. The straight extrapolation of the HE emission is challenged by the VHE upper limit for a dozen objects. For all but two of them, this spectral curvature can be accounted for by the interaction of $\\gamma$~rays with the EBL.  Active galactic nuclei observations, which are crucial both for the understanding of the EBL and of the objects themselves, will remain a primary goal of \\hess\\ during its second phase, \\hess~II, where observations at lower energies will increase the number of detected sources and the maximum redshift accessible by Cherenkov telescopes. Extensive campaigns probing the sky down to fractions of percent of the Crab Nebula flux remain a major task of current VHE telescopes. This tremendous effort is paving the way for targeted AGN observations with CTA."
    },
    "1402/1402.0337_arXiv.txt": {
        "abstract": "Mars Global Surveyor measurements revealed that the Martian crust is strongly magnetized in the southern hemisphere while the northern hemisphere is virtually void of magnetization. Two possible reasons have been suggested for this dichotomy: A once more or less homogeneously magnetization may have been destroyed in the northern hemisphere by, for example, resurfacing or impacts. The alternative theory we further explore here assumes that the dynamo itself produced a hemispherical field \\citep{Stanley2008,Amit2011}. We use numerical dynamo simulations to study under which conditions a spatial variation of the heat flux through the core-mantle boundary (CMB) may yield a strongly hemispherical surface field. We assume that the early Martian dynamo was exclusively driven by secular cooling and we mostly concentrate on a cosine CMB heat flux pattern with a minimum at the north pole, possibly caused by the impacts responsible for the northern lowlands. This pattern consistently triggers a convective mode which is dominated by equatorially anti-symmetric and axisymmetric (EAA, \\citet{Landeau2011}) thermal winds. Convective up- and down-wellings and thus radial magnetic field production then tend to concentrate in the southern hemisphere which is still cooled efficiently while the northern hemisphere remains hot. The dynamo changes from an $\\alpha^2$- for a homogeneous CMB heat flux to an $\\alpha\\Omega$-type in the hemispherical configuration. These dynamos reverse on time scales of about 10 kyrs. This too fast to allow for the more or less unidirectional magnetization of thick crustal layer required to explain the strong magnetization in the southern hemisphere. ",
        "introduction": "Starting in 1998 the space probe Mars Global Surveyor (MGS) delivered vector magnetic field data from orbits between $185$ and $400$ km above the planets surface \\citep{Acuna1999}. The measurements reveal a strong but heterogeneous crustal magnetization \\citep{Acuna1999,Connerney2001}. The more strongly magnetized rocks are mainly localized in the southern hemisphere where the crust is thick and old. The northern hemisphere is covered by a younger and thinner crust which is much weaker magnetized.  Two alternative types of scenarios are discussed to explain this dichotomy. One type explores the possibility that an originally more or less homogeneous magnetization was partly destroyed by resurfacing events after the demise of the internal dynamo. Based on the fact that the Hellas and Argyre impact basins are largely void of magnetization, \\citet{Acuna2001} conclude that the dynamo stopped operating in the early Noachian before the related impact events happened roughly $3.7-4$ Gyrs ago. Volcanic activity and crustal spreading are two other possibilities to explain the lack of strong magnetization in certain surface areas \\citep{Lillis2008b, Mohit2004}, in particular the northern hemisphere after the dynamo cessation.  The alternative scenario explains the dichotomy by an ancient Martian dynamo that inherently produced a hemispherical magnetic field. Numerical dynamo simulations by \\citet{Stanley2008} and \\citet{Amit2011} show that this may happen when more heat is allowed to escape the core through the southern than through the northern core mantle boundary (CMB). Such north/south asymmetry can for example be caused by larger impacts or low-degree mantle convection \\citep{Roberts2006,Keller2009,Yoshida2006}. Due to depth-dependent viscosity and a possible endothermic phase transition \\citep{Harder1996} Martian mantle convection may be ruled in an extreme case by one gigantic plume typically evoked to explain the dominance of the volcanic Tharsis region. However, the single plume convection might have developed after the dynamo ceased. Due to the hotter temperature of the rising material the CMB heat flux can be significantly reduced under such a plume. Though Tharsis is roughly located in the equatorial region it could nevertheless lead to magnetic field with the observed north-south symmetry, as we will show in the following.  The possible effects of large impacts on planets and the dynamo in particular are little understood. \\citet{Roberts2009} argue that impacts locally heat the underlying mantle and thereby lead to variations in the CMB heat flux. Large impacts may also cause a demise of the dynamo by reducing the CMB heat flux below the value where subcritical dynamo action is still possible \\citep{Roberts2009}. The deposition of heat in the outer parts of the core by impact shock waves could lead to a stably stratified core and thereby also stop dynamo action \\citep{Arkani-Hamed2010} for millions of years until the heat has diffused out of the core. If the iron content of the impactor is large enough it may even trigger a dynamo \\citep{Reese2010}.  The thermal state of the ancient Martian core is rather unconstrained \\citep{Breuer2010}. Analysis of Martian meteorites suggests a significant sulphur content and thus a high core melting temperature \\citep{Dreibus1985}. Mars may therefore never have grown a solid inner core, an assumption we also adopt here \\citep{Schubert1990, Breuer2010}. The ancient Martian dynamo was then exclusively driven by secular cooling and radiogenic heating and has stopped operating when the CMB heat flux became subadiabatic \\citep{Stevenson1983}. Run-away solidification or light element saturation may explain the dynamo cessation in the presence an inner core.  The geodynamo, on the other hand, is predominantly driven by the latent heat and light elements emanating from the growing inner core front. Secular cooling and radiogenic heating is typically modeled by homogeneously distributed internal buoyancy sources while the driving associated to inner core growth is modeled by bottom sources \\citep{Kutzner2000}. These latter sources have a higher Carnot efficiency and would likely have kept the Martian dynamo alive if the planet would have formed an inner core.  Several authors have explored the influence of a CMB heat flux pattern on dynamos geared to model Earth and report that they can cause hemispherical variations in the secular variation \\citep{Bloxham2000,Christensen2003,Amit2006}, influence the reversal behavior \\citep{Glatzmaier1999,Kutzner2004} or lead to inhomogeneous inner core growth \\citep{Aubert2008}. However, the effects where never as drastic as those reported by \\citet{Stanley2008} or \\citet{Amit2011}. In the work of \\citet{Amit2011} the reason likely is the increased susceptibility of internally driven dynamos to the thermal CMB boundary condition \\citep{Hori2010}. \\citet{Stanley2008} retain bottom driving and employed a particularly strong heat flux variation to enforce a hemispherical field   Here, we follow \\citet{Amit2011} in exploring the effects of a simple sinusoidal CMB heat flux variation on a dynamo model driven by internal heat sources. The main scope of this paper is to understand the particular dynamo mechanism, to explore its time dependence, and to extrapolate the results to the Martian situation. \\Secref{Model} introduces our model, whereas \\secref{hemisphericalaction} describes the effects of the CMB heat flux anomaly on the convection and the induction process. In \\secref{applymars} we explore the applicability to the ancient Martian dynamo. The paper closes with a discussion in \\secref{Discussion}. ",
        "conclusions": "\\label{Discussion}  We find that an equatorially anti-symmetric convective mode is consistently triggered by a cosine heat flux variation that allows more heat to escape through the southern than through the northern outer boundary of the dynamo region. When the variation is strong enough, convective up- and down-wellings are concentrated at the southern hemisphere and the northern hemisphere remains hot. The associated latitudinal temperature gradients drive strong thermal winds that dominate the flow when, for example, the variation amplitude $g$ exceeds $50\\,$\\% at $E=10^{4}$. Tilting the heat flux pattern axis leaves the solution more or less unchanged with the exception of the $90^\\circ$-case where the equatorial symmetry remain unbroken. We conclude that breaking the equatorial symmetry is dynamically preferred over an equatorially oriented heat flux anomaly of the CMB heat flux.  Due to the thermal winds, the dynamo type changes from $\\alpha^2$ to $\\alpha\\Omega$ but is generally less efficient. Lorentz forces associated with the toroidal field created via the $\\Omega$-effect tend to kill whatever remains of classical columnar convection. This further increases the equatorial anti-symmetry of the solution. Poloidal fields are mainly produced by the southern up- and downwellings which lead to a hemispherical field pattern at the outer boundary.   When the hemisphericity approaches values of that found in Martian crustal magnetization, however, all dynamos start to oscillate on (extrapolated) time scales of the order of $10\\,$kyr. These oscillations are reminiscent of previously described Parker waves in dynamo simulations \\citep{Busse2006}. As a typical characteristic of Parker waves, the frequency increases with the (square root of the) shear strength, see table \\ref{Tab1}. The oscillation periods are much shorter than the time over which the deep reaching Martian magnetization must have been acquired \\citep{Langlais2004}. Being a composite of many consecutive layers with alternating polarities the net magnetization would scale with the time averaged field and would therefore likely be much smaller than the predicted strength of the ancient Martian field magnetizing the crust \\citep{Weiss2002}. The maximum hemisphericity for non-oscillatory dynamos amounts to a configuration where the mean northern field amplitude is only $50\\%$ weaker than the southern. Additional effects like lava-overflows would then be required to explain the observed hemisphericity.  \\citet{Amit2011, Stanley2008} also studied the effects of the identical sinusoidal boundary heat flux pattern and find very similar hemispherical solutions. \\citet{Amit2011} used a very similar setup to ours and also reported oscillations when the dynamo becomes strongly hemispherical. \\citet{Stanley2008} do not report the problematic oscillations intensively studied here, which may have to do with differences in the dynamo models. They study stress-free rather than rigid flow boundaries and assume that the growing inner core contributes to drive the dynamo while our model exclusively relies on internal heating. Should a hemispherical dynamo indeed be required to explain the observed magnetization dichotomy, this may indicate that ancient Mars already had an inner core. Alternatively efficient demagnetization mechanisms may have modified an originally more or less homogeneous magnetized crust \\citep{Shahnas2007}.  \\citet{Landeau2011} observed that similar hemispherical dynamos are found when the Rayleigh number exceeds a critical value. However, albeit the effects are significantly smaller than when triggered via the boundary heat flux.  All the cases explored here remain below this critical Rayleigh number. \\citet{Landeau2011} also mentioned that the equatorial anti-symmetry, and thus the hemisphericity of the magnetic field, decreases when the Ekman number is decreased. Our simulations at lower Ekman number seem to confirm this trend although a meaningful extrapolation to the Martian value of $E=3\\times 10^{-15}$ would require further simulations at lower Ekman numbers. To a certain extent the decrease can be compensated by increasing the heat flux variation amplitude, the Rayleigh number or the magnetic Prandtl number.   Our results show that a north-south symmetry breaking induced by lateral CMB heat flux variations can yield surprisingly strong effects. Fierce thermal winds and local southern upwellings take over from classical columnar convection and the dynamo changes from an $\\alpha^2$ to an $\\alpha \\Omega$-type. The dominant $\\Omega$-effect seems always linked to Parker-wave-like field oscillations typically discussed for stellar applications. It will be interesting to further explore the aspects independent of the application to Mars."
    },
    "1402/1402.4044_arXiv.txt": {
        "abstract": "The classic optical nebular diagnostics [\\ion{N}{2}], [\\ion{O}{2}], [\\ion{O}{3}], [\\ion{S}{2}], [\\ion{S}{3}], and [\\ion{Ar}{3}] are employed to search for evidence of non-Maxwellian electron distributions, namely $\\kappa$~distributions, in a sample of well-observed Galactic \\ion{H}{2} regions. By computing new effective collision strengths for all these systems and $A$-values when necessary (e.g. \\ion{S}{2}), and by comparing with previous collisional and radiative datasets, we have been able to obtain realistic estimates of the electron-temperature dispersion caused by the atomic data, which in most cases are not larger than $\\sim 10$\\%. If the uncertainties due to both observation and atomic data are then taken into account, it is plausible to determine for some nebulae a representative average temperature while in others there are at least two plasma excitation regions. For the latter, it is found that the diagnostic temperature differences in the high-excitation region, e.g. $T_e$(\\ion{O}{3}), $T_e$(\\ion{S}{3}), and $T_e$(\\ion{Ar}{3}), cannot be conciliated by invoking $\\kappa$~distributions. For the low excitation region, it is possible in some, but not all, cases to arrive at a common, lower temperature for [\\ion{N}{2}], [\\ion{O}{2}], and [\\ion{S}{2}] with $\\kappa\\approx 10$, which would then lead to significant abundance enhancements for these ions. An analytic formula is proposed to generate accurate  $\\kappa$-averaged excitation rate coefficients (better than 10\\% for $\\kappa \\geq 5$) from temperature tabulations of the Maxwell--Boltzmann effective collision strengths. ",
        "introduction": "A fundamental and still unresolved problem in the understanding of astronomical photoionized plasmas, particularly \\ion{H}{2} regions and planetary nebulae, is the systematic higher metal abundances derived from recombination lines (RL) relative to those from collisionally excited lines (CEL). The magnitude of these discrepancies varies from a few percent up to puzzling factors as large as 70 \\citep[see, for instance,][]{liu00, liu01}. Several plausible explanations have been proposed to account for their cause, among them temperature fluctuations \\citep{pei67}, hydrogen deficient inclusions \\citep{liu00, sta07}, and dense clumps ionized by X rays \\citep{erc09}. However, all these hypotheses present certain limitations and no conclusive evidence has been put forward to either prove or dispute them.  \\citet{nic12} have suggested that the apparent temperature variations in nebular plasmas are the result of electron velocity distributions that deviate from the generally assumed Maxwell--Boltzmann (MB), proposing the $\\kappa$~distribution as a more functional alternative. It is therein argued that excitation rates sufficiently boosted from the MB values are obtained with $10\\leq \\kappa\\leq 20$, implying that the traditional temperature diagnostics from forbidden line ratios systematically overestimate the true kinetic temperature; therefore, ion abundances derived from CEL would be underestimated. In their analysis, two crucial approximations are made: (i) electron impact collision strengths are constant with energy, which in  practice is not the case due to strong resonance contributions; and (ii) the effects of the $\\kappa$~distribution on the de-excitation rate coefficients are ignored.  \\citet{bin12} have found that the \\ion{O}{3} temperatures from CEL are consistently higher than those of \\ion{S}{3} in an extensive group of both Galactic and extragalatic \\ion{H}{2} regions with gas metallicities higher than 0.2 solar. Among the causes for this excess, the $\\kappa$~distribution has again been invoked. \\citet{nic13} have continued their work on the $\\kappa$~distributions by revising the effects of different datasets of effective collision strengths on the electron temperatures obtained from nebular diagnostics such as [\\ion{N}{2}], [\\ion{O}{2}], [\\ion{O}{3}], [\\ion{S}{2}], and [\\ion{S}{3}]. In particular, by comparing the \\ion{O}{3} electron temperature obtained using recent collision strengths \\citep{pal12} with those derived from earlier datasets, they have concluded that the previous electron temperatures have been seriously overestimated. Since any spectroscopic evidence of non-Maxwellian electron distributions must rely on accurate atomic data, it is indispensable to verify this claim. It must be added that \\citet{nic13} did not consider the accuracy of the radiative rates ($A$-values) which is also a source of uncertainty in spectral diagnostics. For a full treatment of the propagation of atomic data uncertainties in spectral diagnostics, see \\citet{bau13}.  More recently, \\citet{dop13} have made an effort to estimate both MB and $\\kappa$-averaged collisional rates from energy tabulations of raw collision strengths, and stress the need for a systematic inventory and assessment of the newer atomic data and their effects on photoionization models before really establishing the prevalence of the $\\kappa$~distribution in nebular plasmas. This is in fact the main motive of the present work. We make use of the familiar nebular diagnostics---namely [\\ion{N}{2}], [\\ion{O}{2}], [\\ion{O}{3}],  [\\ion{S}{2}], [\\ion{S}{3}], and [\\ion{Ar}{3}]---to search for evidence of $\\kappa$~distributions in well-observed \\ion{H}{2} regions. In Section~\\ref{rates} we describe the effects of the $\\kappa$~distribution on both the electron excitation and de-excitation rate coefficients for forbidden transitions. Since energy-tabulated electron collision strengths with at least a similar degree of reliability are required in this analysis, we decided to recalculate them from scratch (Section~\\ref{upsilon}). The sensitivity of the electron temperature to the $A$-values is also evaluated in Section~\\ref{aval}. With the new data, we derive in Section~\\ref{h2temp} electron temperatures for the \\ion{H}{2} regions studied by \\citet{nic12} and compare them with those obtained with representative effective collision strengths published in the past two decades. The possibilities of the $\\kappa$~distribution in reducing CEL temperatures and in conciliating the temperature differences in plasma regions are also investigated (Section~\\ref{kappaeffects}). Conclusions are finally discussed in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  It has become clear in the work of \\citet{nic13} and \\citet{dop13} on non-Maxwellian electron distributions in nebular plasmas  that reliable effective collision strengths and $A$-values are essential before any recommendations can be drawn. We have therefore been encouraged to make an attempt at quantifying the impact of such atomic data on the electron temperatures of a sample of Galactic \\ion{H}{2} regions for which accurate spectra are available, and to search for evidence to substantiate departures from the MB framework. This initiative has implied extensive computations of new collisional datasets for the ionic systems that give rise in nebulae to CEL plasma diagnostics in order to derive both MB and $\\kappa$-averaged rate coefficients with a comparable degree of accuracy. These energy-tabulated collision strengths are currently available for download from the AtomPy\\footnote{\\tt http://bit.ly/K5oAfD} data curation service.  By extensive comparisons with other datasets of effective collision strengths computed in the past two decades, we have found that their statistical consistency is around the $\\sim 20{-}30$\\% level although larger discrepancies (factor of 2, say) may be detected for the odd transition. This inherent dispersion, which in our opinion would be hard to reduce in practice, is due in forbidden transitions to the strong sensitivity of the resonance contribution to the scattering numerical approach, target atomic model, energy mesh, and small interactions such as relativistic corrections. However, with the exception of the difficult [\\ion{S}{2}] system, the theoretical temperature dispersion in most diagnostics has been found to be not larger than 10\\%, and the agreement with the temperatures estimated in the observational papers is within this similar satisfactory range. Therefore, there is room for optimism. As shown in the present analysis, $T_e$(\\ion{S}{2}) has been proven to be strongly dependent on both the radiative and collisional datasets, and in this context, only the $A$-values by \\citet{men82}, \\citet{fri99}, and the present comply with all the stringent specifications required. On the other hand, the collisional dataset for this ion still needs more refinement before a firm ranking can be put forward on its accuracy.  In the present study, which has relied on high-quality sets of astronomically observed line intensities for \\ion{H}{2} regions with fairly low electron temperatures ($T_e\\lesssim 1.5\\times 10^4$~K), it has been shown that the error margins due to the atomic data are not always taken into account in observational work when quoting electron temperature uncertainties. When they are included, rather than trying to differentiate plasma regions, an average temperature may be defined in some sources with a standard deviation comparable to the error margins of each temperature diagnostic. If multi-region plasmas are indeed in evidence, the observed diagnostic temperature variation in our sample might be resolved by $\\kappa$~distributions only in the low-excitation region, thus leading to a single lower temperature for diagnostics such as [\\ion{N}{2}], [\\ion{O}{2}], and [\\ion{S}{2}] with the consequent abundance enhancements. However, even in these cases the differences between observed line ratios and the predictions of the MB distribution differ by less than $2\\sigma$ if proper uncertainties are taken into account. In the high-excitation region, on the other hand, diagnostic temperature differences in [\\ion{O}{3}], [\\ion{S}{3}], and [\\ion{Ar}{3}] do not seem to be reconciled by invoking $\\kappa$~distributions, which then leaves the problem open to other unaccounted physical processes. It must be emphasized that these findings are not the rule as indicated by the extensive sample by \\citet{bin12} of \\ion{H}{2} regions with mostly higher temperature ($T_e> 10^4$~K); thus further detailed work would be required before $\\kappa$ distributions can be firmly established.  Regarding the effects of the $\\kappa$~distribution on the collisional rate coefficients, we have confirmed the reliability of their neglect for de-excitation, and an accurate (better than 10\\%) analytic expression has been derived to readily estimate $\\kappa$-averaged rates for excitation from temperature tabulations of MB effective collision strengths. The measure that regulates level-excitation departures from MB is $T/T_{\\rm ex}\\ll 1$; hence, for IR transitions within the ionic ground terms with very low excitation energies ($T_{\\rm ex} < 10^3$~K), the MB and $\\kappa$-distribution electron impact excitation rates are essentially identical at nebular temperatures. IR abundance discrepancies with respect to the optical would then be a further indication of the relevance of $\\kappa$~distributions. However, such discrepancies have not been reported; for instance, all planetary nebulae observed with {\\it ISO} show consistent abundances as estimated from both IR and optical CEL \\citep{wes05}, which then suggests that the MB distribution dominates.  In the present work it has been demonstrated that error-margin capping in the collisional data relevant to nebular diagnostic implies---apart from detailed comparisons with previous datasets and in some cases recalculations---comprehensive benchmarks with spectral models and observations, and this is a task we intend to pursue in the ongoing discussion concerning the CEL--RL abundance discrepancies."
    },
    "1402/1402.2809_arXiv.txt": {
        "abstract": "{This paper reports a search for spatial clustering of the arrival directions of high energy muon neutrinos detected by the ANTARES neutrino telescope. An improved two-point correlation method is used to study the autocorrelation of 3058 neutrino candidate events as well as cross-correlations with other classes of astrophysical objects: sources of high energy gamma rays, massive black holes and nearby galaxies. No significant deviations from the isotropic distribution of arrival directions expected from atmospheric backgrounds are observed.}  \\begin{document} ",
        "introduction": "\\label{sec:intro} The key question to resolve the long standing mystery of the origin of cosmic rays is to locate their sources and understand the mechanisms that accelerate these particles up to energies orders of magnitude above the energies reached by man-made accelerators. Over the last years it has become obvious that multiple messengers will be needed to achieve this task. Fundamental particle physics processes like the production and subsequent decay of pions in interactions of high energy particles predict that the acceleration sites of cosmic rays should also be sources of high energy gamma rays and neutrinos. The detection of astrophysical neutrinos and the identification of their sources is one of the main goals of neutrino telescopes operated at the South Pole (IceCube~\\cite{IceCube}), in Lake Baikal~\\cite{Baikal} and in the Mediterranean Sea (ANTARES~\\cite{Antares_DetectorPaper}).  Despite significant effort, no clear signature for point-like sources of astrophysical neutrinos has been found so far~\\cite{Amanda_PointSources2009, IceCube_PointSources2013, IceCube_ICRC2013_PS, ANTARES_PointSources2010, ICRC2013_ANTARESExcess, ICRC2013_PointSources}. Currently, both the spatial distribution as well as the morphologies of sources potentially emitting neutrinos in the TeV energy range are unknown. Similar to the distribution of observed sources emitting high energy gamma rays, they are supposed to be distributed very inhomogeneously throughout our cosmic neighbourhood. A significant fraction of them should be located in the Galactic disk and be spatially extended (e.g. shell-type supernova remnants). It is therefore interesting to study the intrinsic clustering of the arrival directions of neutrino candidates.  In this analysis an improved autocorrelation method is used to this end. As no prior information about the potential sources is required biases are naturally reduced. Since it covers a large angular range, i.e. neutrino emission regions of very different sizes, this study complements the searches for point-like sources and could provide hints for underlying, yet unresolved, source morphologies and source distributions. Exploiting the expected multi-messenger signatures of potential sources the introduced method is extended to searches for correlations between the arrival directions of neutrino candidates and other classes of astrophysical objects: sources of high energy gamma rays, massive black holes and nearby galaxies.  \\subsection{The ANTARES neutrino telescope} \\begin{figure*}[!t] \\centerline { \\includegraphics[width=0.8\\textwidth]{fig01_layout} } \\caption{Schematic view of the ANTARES telescope. The inset shows a photograph of an optical storey.} \\label{fig:layout} \\end{figure*} \\noindent%  The ANTARES telescope~\\cite{Antares_DetectorPaper} became fully operational in 2008. The detector comprises twelve detection lines anchored at a depth of $2475~\\mathrm{m}$ and $40~\\mathrm{km}$ off the French coast near Toulon. The detector lines are about $450~\\mathrm{m}$ long and host a total of 885 optical modules (OMs), each comprising a 17'' glass sphere which houses a 10'' photomultiplier tube. The OMs look downward at $45^\\circ$ in order to optimise the detection of upgoing, i.e. neutrino induced, tracks. The geometry and size of the detector make it sensitive to extraterrestrial neutrinos in the TeV-PeV energy range. A schematic layout of the telescope is shown in Figure~\\ref{fig:layout}.  The neutrino detection is based on the induced emission of Cherenkov light by high energy muons originating from charged current neutrino interactions inside or near the instrumented volume. All detected signals ({\\it hits}) are transmitted via an optical cable to a shore station, where a computer farm filter the data for coincident signals in several adjacent OMs. The muon direction is then determined by maximising a likelihood which compares the time of the hits with the expectation from the Cherenkov signal of a muon track. Details on the event reconstruction are given in Ref.~\\cite{ANTARES_PointSources2010, ANTARES_NuSpectrum}.  Two main backgrounds for the search for astrophysical neutrinos can be identified: downgoing atmospheric muons which have been mis-reconstructed as upgoing and atmospheric neutrinos originating in cosmic ray induced air showers at the opposite side of the Earth. Depending on the requirements of the analysis both backgrounds can, at least partially, be discriminated using various parameters such as the quality of the event reconstruction or an estimator of the energy of the muon, e.g. the number of hits used in the track reconstruction. The latter, being strongly correlated with the energy of the original neutrino, helps to discriminate events of atmospheric origin from neutrinos produced in astrophysical sources. Atmospheric neutrinos have a much softer energy spectrum ($\\propto E^{-3.7}$) compared to the generic $E^{-2}$ spectrum expected from Fermi acceleration in astrophysical sources. As illustrated in Figure~\\ref{fig:nHit} this difference affects the distribution of energy dependent parameters (or {\\it energy estimators}) and can therefore be used to enhance the background discrimination. In addition, analysing the reconstructed arrival directions of the events allows to search for an excess over the isotropic atmospheric backgrounds. Both features will be exploited in the analysis described in this paper.  \\begin{figure*}[!t] \\centerline { \\includegraphics[width=0.75\\textwidth]{fig02_nHits} } \\caption{Normalised distribution of the number of hits used in the event reconstruction for data (black markers) and Monte Carlo simulations (atmospheric neutrinos following the parametrisation from Ref.~\\cite{Bartol}: black histogram; astrophysical neutrinos: red, dotted histogram).} \\label{fig:nHit} \\end{figure*} \\noindent  \\subsection{Dataset} The data analysed here has been recorded by the ANTARES neutrino telescope between 2007 and 2010. During the beginning of this period (2007-2008) the detector was in its commissioning phase, increasing from five active detection lines to the full detector with twelve lines by mid 2008. After imposing basic data quality requirements, the event selection criteria have been optimised by means of Monte Carlo simulations to yield the best average upper limit on the neutrino flux in a search for point-like sources~\\cite{ANTARES_PointSources2010}. These criteria are mainly a cut on the reconstructed zenith angle, $\\theta > 90^\\circ$, a requirement on the reconstruction quality parameter, $\\Lambda$, as well as a cut on the estimated angular uncertainty of the track reconstruction, $\\beta < 1^\\circ$. A total of 3058 neutrino candidates are found in 813 days of effective lifetime. Monte Carlo simulations show that the contribution from misreconstructed atmospheric muons is about $15~\\%$ in this dataset. A skymap in galactic coordinates of these events  is shown in the upper left plot of Figure~\\ref{fig:SkyMaps}. ",
        "conclusions": "In the search for the sources of high-energy cosmic rays, the detection of astrophysical neutrino sources may play a crucial role. Here a search for intrinsic clustering of data recorded with the ANTARES neutrino telescope is presented. This analysis uses an improved two-point correlation technique exploiting an estimate of the energy of the neutrino candidates. The arrival directions of the selected neutrino candidates neither show evidence for clustering of events on top of the isotropic distribution expected for the background of atmospheric neutrinos, nor correlate with catalogues of gamma rays, nearby galaxies or massive black holes."
    },
    "1402/1402.2512_arXiv.txt": {
        "abstract": "We present a family of self-consistent axisymmetric stellar systems that have analytic distribution functions (\\df s) of the form $f(\\vJ)$, so they depend on three integrals of motion and have triaxial velocity ellipsoids. The models, which are generalisations of H\\'enon's isochrone sphere, have four dimensionless parameters, two determining the part of \\df\\ that is even in $L_z$, and two determining the odd part of the \\df\\ (which determines the azimuthal velocity distribution). Outside their cores, the velocity ellipsoids of all models tend to point to the model's centre, and we argue that this behaviour is generic, so near the symmetry axis of a flattened model, the long axis of the velocity ellipsoid is naturally aligned with the symmetry axis and not perpendicular to it as in many published dynamical models of well-studied galaxies.  By varying one of the \\df's parameters, the intensity of rotation can be increased from zero up to a maximum value set by the requirement that the \\df\\ be non-negative. Since angle-action coordinates are easily computed for these models, they are ideally suited for perturbative treatments and stability analysis.  They can also be used to choose initial conditions for an N-body model that starts in perfect equilibrium and to model observations of early-type galaxies.  The modelling technique introduced here is readily extended to different radial density profiles, more complex kinematics, and multi-component systems. A number of important technical issues surrounding the determination of the models' observable properties are explained in two appendices. ",
        "introduction": "Although real galaxies are by no means in states of dynamical equilibrium, equilibrium models have a fundamental role to play in the study of galaxies and the Universe. One reason for this is that most of a galaxy's mass is thought to reside in dark matter that can only be traced by its gravitational field, and we can map this field through the kinematics of field stars only to the extent that the stars are in dynamical equilibrium. A further reason is that modest deviations of galaxies from dynamical equilibrium are best modelled by perturbing equilibrium models -- most successful branches of physics, from celestial mechanics to high-energy physics through plasma physics and quantum condensed-matter physics, comprise applications of perturbation theory.  Currently the range of equilibrium galaxy models is extraordinarily limited. The theory of globular clusters rests to a great extent on the equilibrium models introduced by \\cite{Michie} and popularised by \\cite{King}. Our still very limited understanding of spiral structure owes much to the two-dimensional equilibrium models of \\cite{Kalnajs}, Zang \\& Toomre \\citep{Toomre} and Evans \\citep{Evans94,ReadEvans,EvansSellwood}.  On account of the paucity of fully three-dimensional equilibrium models, several recent analyses of our Galaxy's halo have relied on  models which have distribution functions (\\df s) of the form \\[\\label{eq:beta} f(E,L)=L^{-2\\beta}F(E), \\] where $E$ is a star's energy and $L$ is the magnitude of its angular momentum \\citep{Deasonxx}. Models with a \\df\\ of this form only make dynamical sense to the extent that the gravitational potential can be considered to be spherically symmetric, which it probably cannot be in the solar neighbourhood. Moreover, in the case of radial bias, the \\df\\ (\\ref{eq:beta}) implies infinite phase-space density in the limit $L\\to0$ of radial orbits, which is inherently implausible and potentially compromises the model's stability \\citep{FridP,PalmerP}, while in the case $\\beta<0$ of tangential bias, the \\df\\ implies a distribution of velocities that is physically implausible because it is bimodal in $v_\\phi$ \\citep{FermaniSa}. Hence \\df s of the form (\\ref{eq:beta}) do not constitute a satisfactory basis for dynamical modelling.  Although galaxies are probably never precisely axisymmetric, they are often sufficiently nearly so for axisymmetric models to be valuable starting points from which better-fitting models may be derived by perturbation theory. In this paper we present a new class of axisymmetric models. In a subsequent paper it will be shown how the technique we introduce here can be used to produce a wide range of axisymmetric models.  Already in the paper in which Jeans introduced his theorem, it was clear that the equilibrium \\df\\ of our Galaxy cannot be a function of the form $f(E,L_z)$ that depends only on the classical integrals in an axisymmetric potential. Dependence on a third integral $I_3$ is essential, and the field has been held back by the want of an analytic formula for $I_3(\\vx,\\vv)$. In the 1970s it became evident that the \\df s of elliptical galaxies also involve $I_3$ in an essential way \\citep{Bertola,Binney76,Daviesetal}, although recent systematic surveys have cast a new light on this conclusion \\citep{Cappellarietal,Emsellem}. Nevertheless, it remains true that an abundance of observational material indicates that realistic equilibrium models of real galaxies must depend on $I_3$ in addition to $E$ and $L_z$.  Numerical orbit integrations in the 1960s showed that generic orbits in flattened axisymmetric potentials respect a third integral $I_3$ \\citep{HenonH,Ollongren}. But these experiments did not lead to useful analytic expressions for $I_3(\\vx,\\vv)$. A promising attack on this problem through Hamiltonian perturbation theory was pursued by Contopoulos and his collaborators \\citep{Contopoulos}, but this line of attack was frustrated by the fact that the vertical oscillations of most stars are far from harmonic, so their orbits are not readily treated as perturbed harmonic oscillators. Moreover the coupling between a star's radial and vertical oscillations is fundamental to their dynamics, so neither motion should be considered in isolation.  \\cite{Eddington} and \\cite{Stackel} showed that analytic expressions for $I_3$ can be obtained for a certain class of potentials that are now known as St\\\"ackel potentials.  \\cite{deZeeuw} showed that many of these potentials are generated by remarkably galaxy-like density distributions, and he clarified the nature of orbits in these potentials. The present paper relies on a technique, the ``St\\\"ackel Fudge'' \\citep[][hereafter B12a]{Binney12a}, which is an extension of this classic work. This approximation consists of applying to an arbitrary gravitational potential $\\Phi(R,z)$ formulae that would be valid if the potential were of St\\\"ackel's form even though the potential does not have this form.  \\begin{figure} \\centerline{\\epsfig{file=tetra.ps,width=.6\\hsize}} \\caption{The tetrahedral zone in action space occupied by orbits with energies less than some maximum value.} \\label{fig:Jspace} \\end{figure}  Since any function of a star's isolating integrals is itself an isolating integral, there is in principle  enormous freedom in the choice of arguments of a galaxy's \\df. However, certain integrals stand out for special consideration: the action integrals. These alone can be embedded in a canonical coordinate system, and their conjugate (angle) variables have the remarkable property of increasing linearly in time: \\[\\label{eq:theta_of_t} \\theta_i(t)=\\theta_i(0)+\\Omega_i t. \\] Action integrals are unique up to a set of discrete canonical transformations that map between rational linear combinations of any given set of actions, and map the angle variables into integer linear combinations of the given angles. Hence when combined with reasonably physical requirements such as ``the radial action $J_r$ should quantify the extent of radial oscillations'' and ``the vertical action $J_z$ should quantify the extent of vertical oscillations'', actions are uniquely defined. This uniqueness greatly facilitates the comparison of models by making it possible to compare the density with which orbits in slightly different potentials are populated.  Therefore it is natural to require the \\df\\ to be a function $f(\\vJ)$ of the actions.  To derive from $f(\\vJ)$ the observable properties of a model, we need expressions for the actions as function $\\vJ(\\vx,\\vv)$ of the ordinary phase-space coordinates. Here we use the St\\\"ackel Fudge from B12a.  Numerical experiments presented in B12a show that the actions and angles that one obtains from the Fudge are significantly more accurate than those obtained from the adiabatic approximation \\citep{B10}, and they are also valid for orbits that move far from the equatorial plane and thus lie outside the range of validity of the adiabatic approximation.  Whereas \\cite{Binney12b} and \\cite{Binneyetal13} only required actions for orbits that pass within $\\sim2\\kpc$ of the solar neighbourhood, we require actions throughout our models. This change has prompted us to undertake significant revisions of the scheme described by B12a to obtain them with increased accuracy and reduced numerical effort. These revisions are described in Appendix I.  In Section \\ref{sec:DF} we define our \\df s, which are derived from the \\df\\ of the isochrone sphere. In Section \\ref{sec:relax} we explain how a model is relaxed to its self-consistent configuration and detail checks of the numerical accuracy.  In Section \\ref{sec:obs} we describe the observable properties of some specific models and refine our choice of \\df. In Section \\ref{sec:appls} we discuss some potential applications of these models, including choice of initial conditions for N-body models, studies of symmetry breaking in stellar systems and related perturbation analyses, and modelling observations of early-type galaxies. Section \\ref{sec:conclude} sums up. Two Appendices give numerical details that are likely to be useful when building the models. ",
        "conclusions": "\\label{sec:conclude}  We have presented a new type of self-consistent model of hot, axisymmetric stellar systems that have specified \\df s $f(\\vJ)$, where $\\vJ$ is the triple of action values. The even part of the \\df\\ is specified by two dimensionless parameters $\\alpha_{\\phi0}$ and $\\alpha_z$. When $\\alpha_z>1$ the model becomes oblate. If $\\alpha_{\\phi0}\\simeq1$ the model is radially biased, while dropping $\\alpha_{\\phi0}$ below unity reduces the radial bias and, for sufficiently small values of $\\alpha_{\\phi0}$, the model becomes azimuthally biased. The odd part of the \\df, which does not contribute to the model's density profile $\\rho(\\vx)$, is controlled by two dimensionless parameters, $\\chi\\ge0$, which controls the steepness of the rotation curve at the centre, and $0\\le k\\le0.5$, which determines how fast the model rotates: increasing $k$ both speeds up the model's rotation and diminishes the magnitude of the azimuthal velocity dispersion $\\sigma_\\phi$.  We focused on the observables of just two exemplary models, a tangentially biased model with $\\alpha_{\\phi0}=0.7$ and a radially biased model with $\\alpha_{\\phi0}=1$.  Both models achieve peak ellipticities $\\epsilon\\simeq0.3$ at about $R_\\e/2$ and are distinctly boxy when seen edge-on.  In each model the velocity ellipsoids are aligned with cylindrical coordinates within the core, but beyond the core they quickly align with radial polar coordinates, and along the minor axis the ellipsoids point towards the centre rather than tangentially. We have argued that the only way to make the velocity ellipsoids point tangentially at significant distances down the minor axis is to include a distinct component of stars on nearly polar orbits.  The models could provide initial conditions for N-body models that start from perfect equilibrium, something that is possible only when the \\df\\ is explicitly known. Since the angle-action coordinates of any point in the phase space of one of these models are readily computed, the models provide perfect testbeds for studies of galactic stability and evolution.  The present models differ from Schwarzschild models in having vastly fewer free parameters: four versus the number of orbits in the orbit library employed. They differ from models based on the Jeans equations in providing full velocity distributions rather than just the first two moments of the distributions. It would be straightforward to extend these models in various directions. For example, it is easy to devise other approaches than simple action scaling to move from the \\df\\ $f(H)$ of an ergodic model to a three-integral \\df\\ f(\\vJ), and we are currently exploring one of these alternatives.  Another direction in which the present work is being extended is to multi-component systems, in which stars and dark matter have distinct distribution functions, and, in the case of our Galaxy, the thin disc, thick disc and halo stellar populations all have distinct \\df s.  We started from the isochrone sphere because it provides an analytic expression for $H(\\vJ)$. In a forthcoming paper  we will show how to obtain a good approximation to an analogous expression for any spherical model, and thus extend the present work to other popular spherical systems, such as the \\cite{Hernquist} sphere."
    },
    "1402/1402.5104_arXiv.txt": {
        "abstract": "The Friedmann-Robertson-Walker (FRW) universe and Bianchi I,II universes are investigated in the framework of the generalized uncertainty principle (GUP) with a linear and a quadratic term in Planck length and momentum, which predicts minimum measurable length as well as maximum measurable momentum. We get a dynamic cosmological bounce for the FRW universe. With Bianchi universe, we found that the universe may be still isotropic  by implementing GUP. Moreover, the wall velocity appears to be stationary with respect to the universe velocity which means that when the momentum of the Universe evolves into a maximum measurable energy, the bounce is enhanced against the wall which means no maximum limit angle is manifested anymore. ",
        "introduction": "The existence of space-like singularities is a generic feature of Einstein's general theory of relativity and this is proved by the singularity theorems of Penrose, Hawking and Geroch\\cite{hawell}. The singularities are usually characterized by divergences of curvature invariants or breakdown of geodesics. This indicates a limit beyond which general theory relativity is not applicable anymore. It is widely believed that a theory of quantum gravity will be the one which can provide insights on the resolution of singularities. In fact there are different approaches for quantum theory of gravity, so the description at low energy limit also have competing candidates even at the phenomenological level.\\par Various approaches for quantum gravity such as string theory and black hole physics, have predicted the existence of a minimum measurable length and an essential modification of the Heisenberg uncertainty principle \\cite{guppapers, BHGUP, kmm, kempf, brau, Scardigli} to the so-called generalized uncertainty principle (GUP). The GUP is based on the modification of the fundamental commutation relation mainly in position and momenta. Recent developments suggests that the minimal length scale can be implemented in quantum mechanics and quantum field theory through the GUP. Within the framework of the GUP  minimal length scale has been studied in details in quantum mechanics, quantum electrodynamics, thermodynamics, black-hole physics and cosmology. For some recent reviews on the phenomenology of different minimal length scale scenarios inspired by various approaches to the quantum gravity see \\cite{HossenNico}. \\par Quite recently Ali et al. \\cite{advplb, Das:2010zf} introduced a new approach which predicts maximum observable momenta besides a minimal measurable length. This model is built to be consistent with DSR\\cite{cg}, string theory and black hole physics \\cite{guppapers, BHGUP}. Furthermore, it ensures that $[x_i,x_j]=0=[p_i,p_j]$ through Jacobi identity\\cite{Ali:2011fa}. Accordingly, the modification of the Heisenberg uncertainty relation near the Planck scale reads  \\bea [x_i,p_j] &=& i \\hbar \\left[\\delta_{i j} - \\alpha \\left(p \\delta_{i j} + \\frac{p_i p_j}{p}\\right) +  \\alpha^2 \\left(p^2 \\delta_{i j} + 3 p_i p_j\\right) \\right]~. \\label{comm01} \\eea  Here $\\alpha$ is the deformation parameter given by $\\alpha=\\alpha_0/(c\\, M_{pl}) =\\alpha_0 l_{pl}/\\hbar$. $c$, $M_{pl}$ and $l_{pl}$ are speed of light in vacuum, Planck mass and Planck length respectively. $\\alpha_0$ is a dimensionless parameter usually considered to be of order unity. The upper bounds on the parameter $\\alpha_0$ have been estimated in \\cite{Ali:2011fa} and it was suggested that it could predict an intermediate length scale between Planck scale and electroweak scale. A recent proposal suggested that these bounds can be measured using quantum optics techniques in \\cite{Nature,NatureGRW} which may be considered as a milestone in quantum gravity phenomenology. Due to the GUP as proposed in \\cite{advplb} the physical momentum is redefined. As a result the classical Hamiltonian as well as the quantum Hamiltonian gets modified which affects the quantum phenomena. Recently, Bekenstein \\cite{Bekenstein:2012yy,Bekenstein:2013ih} proposed that quantum gravitational effects could be tested experimentally, he suggests ``a tabletop experiment which, given state of the art ultrahigh vacuum and cryogenic technology, could already be sensitive enough to detect Planck scale signals'' \\cite{Bekenstein:2012yy}. In a series of papers, the effects of GUP on atomic systems, condensed matter systems, preheating phase of the universe, inflationary era of the universe and black holes have been  investigated \\cite{dvprl,dvcjp, Ali:2011fa,Chemissany:2011nq, Tawfik:2012he, Ali:2012mt, faragali, neutrinoGUP, Majumder:2012qy, Majumder:2012ph, Nozari:2012nf, Nozari:2012gd, majahep, majgrg, majahep2}.\\par  In this paragraph we shed some light on the phase space formulation of quantum mechanics as in the later sections we are going to use GUP modified classical Poisson brackets. The known correspondence between commutator and Poisson bracket was first proposed by Dirac where he proposed that the quantum counterparts $\\hat{A}$, $\\hat{B}$ of classical observables $A$, $B$ satisfy \\cite{1} \\bea [\\hat{A}, \\hat{B}]= i \\hbar \\{A,B\\} \\eea and since that time, this is known as a postulate of Quantum Mechanics. In 1946, Hip Groenewold demonstrated that a general systematic correspondence between quantum commutators and Poisson brackets could not hold consistently \\cite{2}.  However, he did appreciate that such a systematic correspondence does, in fact, exist between the quantum commutator  and a deformation of the Poisson bracket,  which is called the Moyal bracket \\cite{3}. The Moyal bracket is a way of describing the commutator of observables in the phase space formulation of quantum mechanics when these observables are described as functions on phase space. It relies on schemes for identifying functions on phase space with quantum observables, the most famous of these schemes being  Weyl quantization \\cite{4}. It underlies Moyal's dynamical equation, an equivalent formulation of Heisenberg's quantum equation of motion, thereby providing the quantum generalization of Hamilton's equations. In a two-dimensional flat phase space, and for the Weyl-map correspondence, the Moyal  bracket reads, \\bea \\{\\{f,g\\}\\} = \\frac{1}{i\\hbar} (f\\star g- g\\star f)= \\{f,g\\}+\\mathcal{O}(\\hbar^2). \\eea Where $\\star$  is the star-product operator in phase space (Moyal product), while $f$ and $g$ are differentiable phase-space functions, and $\\{f,g\\}$ is their Poisson bracket. This means Moyel bracket is equal to poisson bracket in its equivalence with the quantum commutator up to the second order of $\\hbar$ (i.e $\\hbar^2$). In this present paper we study the implications of GUP which has a linear term in $\\alpha$ as well as a quadratic one. For the Friedmann model we only consider the correction which is proportional to $\\alpha$ in the commutator of eqn.(\\ref{comm01}) and $\\hbar \\alpha\\sim \\alpha_0 \\ell_P\\sim\\hbar^{1/2}$ where $\\ell_P=\\sqrt{\\hbar G/c^3}$, which means Moyel bracket is not necessary in our model and Poisson bracket is quite enough to study the commutator of GUP. For studying Bianchi models we consider the most general form of GUP (including linear as well as quadratic term in $\\alpha$) and also GUP with only linear term in $\\alpha$ in which we trust using Poisson bracket, and  we make a comparative analysis with earlier results in the literature \\cite{montani}.  \\par Motivated by the GUP, many authors studied some classical problems with deformed Poisson brackets \\cite{c36,c38}. From the phenomenological point of view, constraints were placed on deformation parameters in \\cite{c36} by considering the effects of GUP motivated deformed Poisson algebra on the classical orbits of particles in the central force problem. In \\cite{c38} it was conjectured that modified commutation relations as suggested by candidate theories of quantum gravity, persist in the classical limit also. There the perihelion precession rate for Keplerian orbits was calculated. A deformed Heisenberg algebra or in the classical case, a deformed Poisson algebra incorporate some additional problems like the violation of equivalence principle \\cite{faragali}. But recently in \\cite{c45} it was shown that the GUP is reconciled with the equivalence principle. The effects of the GUP on Galilean and Lorentz transformation is also studied recently in \\cite{cmain}.\\par In this paper we are going to study the implication of the GUP as proposed in \\cite{advplb} in early universe cosmology. The effect of $\\alpha$ in (\\ref{comm01}) is relevant in the very early universe or near Planck scale and it dies down at low energy scales and hence can be neglected at a later time in the evolution of the universe. However there are attempts to address the problem of Dark Energy in the realm of GUP. In \\cite{barunMUP} the classical and quantum effects of this GUP are investigated on the phase space of a dilatonic cosmological model and it was found that it is possible to get a late time acceleration for this model. A similar study was done in \\cite{darabi} and it was shown that GUP at first generates acceleration but prevents the eternal acceleration at late times and turns it into deceleration. The GUP can also incorporate corrections in the entropy area relation and thereby modify the energy densities in Holographic and Agegraphic Dark Energy Models. Due to this we can get distinct terms in the form of $f(R)$ which may have its importance in explaining the early inflationary scenario as well as the present accelerated expansion \\cite{barunfr}. \\par The scope of the present work is to investigate the effect of the GUP on the dynamics of the Friedmann-Robertson-Walker universe (FRW) by performing this modification at the classical level by studying the modifications induced on the symplectic geometry by the deformed algebra. It is found that big bang singularity seems to be suppressed, by considering GUP. It is found that  modified Friedmann equation predicts a cosmological bounce. Besides, we extend our study to investigate the effect of GUP on Bianchi I and II universes. In each case the deformed commutation relations due to GUP would modify the dynamical equations and hence we will find departures from the usual scenario. We investigate the isotropy of Bianchi universes as well as the singularity. Later we will see that the results remain qualitatively same compared to the GUP which has a quadratic term in $\\alpha$ \\cite{battisti}. However a discussion on the sign of $\\alpha$ is made later which is important for the analysis. An outline of this paper is as follows: in Sec. \\ref{sec:FRW}, we investigate the GUP with FRW universe and derive the modified Freidmann equation. In Sec. \\ref{Bianchi}, we tackle the impact of GUP on Bianchi universes with type I and II. We compare our results with previous studies in \\cite{montani}. ",
        "conclusions": "In this work we have investigated the consequences of the GUP, which predicts maximum observable momenta besides a minimal measurable length \\cite{advplb,Das:2010zf} and is consistent with DSR\\cite{cg}, string theory and black hole physics \\cite{guppapers, BHGUP}, on FRW universe and Bianchi I and II universes. By investigating the effect of GUP with FRW universe, we found that the big bang singularity seems to be suppressed by a dynamical critical density (dynamical in the sense that it is scale factor dependent). Moreover, we extend our study for Bianchi I and II universes with different forms of GUP available in the literature. By investigating different forms of GUP with Bianchi I, we found that the Universe may possibly be isotropic and may evolve into a situation at which the Kasner indices become equal with implementing GUP. With Bianchi II, we found that  in the asymptotic limit of the highly deformed case ($\\mu>>1$) the maximum angle is given by $\\vert \\theta_i \\vert < \\vert \\theta_{max} \\vert =\\cos^{-1} (\\dot{x}_w/\\dot{x}) = \\pi/2$. when the momentum of the particle (Universe) evolves into a maximum measurable energy, the bounce is enhanced against the wall which means no maximum limit angle is manifested anymore. For low $\\mu$,  we have a maximum of $\\dot{x}_w/\\dot{x}$ for low values of $\\mu$ at around $0.67$ where the phenomenon of bounce is accelerated. The maximum angle for this case is $\\vert \\theta_{max}\\vert \\approx \\frac{4\\pi}{15}$. We conclude that the proposed GUP in this work might possibly resolve singularity problems with the considered universes and may imply a bounce picture for the universe.\\par We should also point out that an expanding universe becomes isotropic due to the matter contributions. But here we have not considered any particular matter which is a shortcoming of our present model. So it is indeed necessary to study the effective equations of our model with different matter contributions which we plan to consider in near future."
    },
    "1402/1402.4484_arXiv.txt": {
        "abstract": "We present the results of an X-ray follow-up campaign targeting 10 Weak Lensing (WL) selected galaxy clusters from a {\\em Subaru} WL survey.  Eight clusters were studied with dedicated {\\em Chandra} pointings, whereas archival X-ray data were used for the remaining two clusters.  The WL clusters appear to fit the same scaling relation between X-ray luminosity and temperature as X-ray selected clusters.  However, when we consider the luminosity-mass relation, the WL selected clusters appear underluminous by a factor 3.9$\\pm$0.9 (or, alternatively, more massive by 2.9$\\pm$0.2), compared to X-ray selected clusters.  Only by considering various observational effects that could potentially bias WL masses, can this difference be reconciled. We used X-ray imaging data to quantify the dynamical state of the clusters and found that one of the clusters appears dynamically relaxed, and two of the clusters host a cool core, consistent with Sunyaev-Zel'dovich effect selected clusters.  Our results suggest that regular, cool core clusters may be over-represented in X-ray selected samples. ",
        "introduction": "\\label{sec:wlintro}  Clusters of galaxies are the largest gravitationally-collapsed structures in the Universe. The study of their number density and growth from high density perturbations in the early Universe offers insight into the underlying cosmology \\citep[e.g.,][]{2008MNRAS.387.1179M,2009ApJ...692.1060V}. However, the use of clusters as a cosmological probe requires an efficient method to find clusters over a wide redshift range, an observational method of determining the cluster mass and a method to compute the selection function or the survey volume in which clusters are found \\citep{2008MNRAS.387.1179M}. A variety of observational techniques has been brought to bear on these  requirements, each with different strengths and weaknesses.  Early samples of clusters were based upon optical selection, however these samples have traditionally suffered from projection effects and uncertainties in the optical richness-mass relation.  Recently, the first clusters have been detected in blind surveys using the Sunyaev-Zel'dovich effect \\citep[SZE;][]{2009ApJ...701...32S,2010ApJ...722.1180V,2011A&A...536A...8P}. This method holds much promise due to the redshift independence of the SZE \\citep{1999PhR...310...97B}, however the use of the SZE as a mass estimator remains largely untested. Until now, the most effective method of building large, well defined cluster samples has been via X-ray selection. The high X-ray luminosities of clusters makes it relatively easy to detect and study clusters out to high redshifts, and X-ray cluster studies have provided a means to impose tight constraints on cosmological parameters \\citep[e.g.][]{2009ApJ...692.1060V,2010MNRAS.406.1759M}.  A common weakness of the three techniques outlined above is that the clusters are selected based on the properties of their minority baryonic content. However, cosmological constraints from those clusters are based on their masses, dominated by dark matter. This gives rise to complications both in the estimation of cluster masses (which must be inferred from the observed baryon properties) and in the determination of sample selection functions. Incorporating selection functions in any cosmological model depends crucially on the form and scatter of the relationship between the observable used to detect the cluster, and the cluster mass.  In principle, these complications may be avoided by defining cluster samples through gravitational lensing, the most direct observational probe of cluster masses. The development of weak lensing (abbreviated WL throughout) techniques has enabled the detection of clusters via the distortions of background galaxy shapes leading to the construction of WL selected cluster surveys. These include the Deep Lens Survey \\citep{2006ApJ...643..128W} and our {\\em Subaru} Weak Lensing Survey \\citep{2007ApJ...669..714M}, both $\\approx$20 deg$^2$, and the large ($\\sim$170 deg$^{2}$) CFHT Legacy Survey \\citep{2012ApJ...748...56S}. However, while WL techniques are free from assumptions about the relationship between baryonic and dark matter in clusters, they are susceptible to the possibility of projection of multiple structures along the line of sight, leading to overestimates of cluster masses or false detections.  This paper aims to determine the X-ray properties of clusters detected via their WL signal. By constructing scaling relations based on the measured X-ray properties, we will determine whether the clusters follow simple scaling theory, which is a key ingredient in the determination of cosmological parameters.  One of the most important measurements for use in cosmological studies is that of the cluster' mass.  Much work has been done to determine cluster masses from weak lensing observations, and the comparison to X-ray observables and X-ray determined masses, however this is primarily based on X-ray selected cluster samples \\citep[e.g.][]{2007MNRAS.379..317H,2010ApJ...721..875O,2011ApJ...737...59J,2012A&A...546A..79I,2012MNRAS.427.1298H,2013ApJ...767..116M,2014arXiv1402.3267I}. We will compare the scaling relation between the X-ray luminosity and WL mass from WL selected and X-ray selected samples. The derivation of the WL cluster mass assumes that the clusters follow spherical symmetry.  We investigate this assumption and determine the dynamical state and cool-core fraction of WL selected clusters.  The outline of this paper is as follows.  In $\\S$~\\ref{sec:wlreduc} we discuss the sample selection and data reduction.  The derivation of the cluster properties is given in $\\S$~\\ref{sec:wlprops}. In $\\S$~\\ref{sec:wlresults} we present the results of our X-ray analysis and derive scaling relations.  Our discussions and conclusions are presented in $\\S$~\\ref{sec:wldisc} and $\\S$~\\ref{sec:wlconclusion} respectively.  Throughout this paper we assume a cosmology with $\\Omega_{\\rm M}$=0.3, $\\Omega_{\\Lambda}$=0.7 and $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:wlconclusion}  We have presented the results of X-ray observations of 10 WL selected galaxy clusters in order to determine their X-ray properties.  Our main results are as follows. \\begin{enumerate} \\item We find that the scaling relation between X-ray luminosity and temperature is consistent with samples of X-ray selected clusters. \\item We find that the WL selected clusters are offset from X-ray selected clusters and groups in the luminosity-temperature relation, implying either the X-ray luminosity is under-estimated or the WL mass is over-estimated. \\item The centroid shifts for the WL clusters show that the majority appears dynamically unrelaxed compared to X-ray selected samples, suggesting X-ray cluster surveys preferentially detect clusters that are morphologically relaxed and/or WL selection favours morphologically disturbed clusters. \\item The cuspiness of the gas density profiles shows that two of the clusters in our sample appear to host a cool core. \\item Measuring the ellipticities of the BCGs for each cluster, and assuming that they are intrinsically prolate, we find that 60$\\%$ of the clusters appear to be viewed with their major axis close to the line of sight. \\item Through a series of corrections to the WL cluster mass including the centroid position, Eddington bias, and triaxiality, and calculating the WL masses with an X-ray derived r$_{500}$, we find the WL selected clusters are consistent in the LM relation with X-ray selected clusters. \\item Comparisons to X-ray selected samples with WL derived cluster masses, and within WL defined r$_{500}$s, show that the offset in the LM relation is still present, even when correcting for centroid position, Eddington bias, and triaxiality. \\end{enumerate}  Our results show that WL selected clusters are affected by biases which, when combined, has a large effect on the calculated WL mass. We have presented a series of possible biases to explain the large offset of the LM relation when compared to X-ray selected cluster samples.  When considering the cumulative effect of these biases on the WL mass we reconcile the difference in the LM relation.  Our comparison to an X-ray selected sample of clusters with WL derived cluster masses, show that even after taking into account biases which may arise due to the WL selection, there is still a discrepancy in the LM relation.  This implies that there maybe further biases in the WL selection, or that the WL selected clusters are embedded in filamentary structure viewed along the line of sight."
    },
    "1402/1402.6953_arXiv.txt": {
        "abstract": "{ There is a great interest in measuring the non-electronic component of neutrinos from core collapse supernovae by observing, for the first time, also neutral-current reactions. In order to assess the physics potential of the ultra-pure  scintillators in this respect, we study  the entire expected  energy spectrum in the Borexino, KamLAND and SNO+ detectors. We examine the various sources of uncertainties in the expectations, and in particular, those due to specific detector features and to the relevant cross sections. We discuss the possibility to identify the different neutrino flavors, and we quantify the effect of confusion, due to other components of the energy spectrum, overlapped with the neutral-current reactions of interest.}    \\begin{document} ",
        "introduction": "\\label{sec:intro}  A Core Collapse Supernova (SN) releases 99\\% of its total energy by emitting neutrinos of the six flavors. The capability to observe the electronic antineutrino component of this emission has already been proven by the detection of SN1987A neutrinos \\cite{KII,IMB,Baksan}. The very large statistics that we will collect from the next galactic supernova  will allow us to study the time dependence of the spectrum, specific features of the $\\bar{\\nu}_e$ luminosity and of its average energy \\cite{SK, IceCube}. The detection of the other neutrinos flavors, however, requires specific detectors and interactions, typically with smaller cross sections. This is true, in particular, for the non electronic component of the spectrum, that can be observed only through Neutral Current (NC) interactions.    During the last years a new generation of ultra-pure liquid scintillators, Borexino (BRX) \\cite{2009BorexinoColl} and KamLAND (KAM) \\cite{2003KamlandColl},  have been operated, obtaining excellent results thanks to the unprecedented low background levels reached and the new sensitivity in the very low energy range, below 1 MeV. They have a particularly good physics potential for the detection supernova NC channels, and quite remarkably, the Elastic Scattering (ES) of (anti)neutrinos on protons \\cite{beato}. It has been argued that the high statistics from this reaction should suffice to constrain the spectra of  the non electronic component for a SN emission, already with the existing detectors  \\cite{dasg}. In view of the importance of this conclusion, we would like to reconsider it in this work.  The outline of this paper is as follows. First of all, we summarize the available information regarding SN neutrino detection in the existing ultra-pure scintillators, and calculate for each of them the total number of expected events as well as their spectral features. We consider the contributions of all neutrino interaction channels and obtain in this way the spectrum of events for a galactic supernova. In this way, we are in the position to evaluate which are the capabilities of the present generation of  ultra-pure scintillators to identify and measure the different neutrino flavors. ",
        "conclusions": "In this paper, we have obtained and discussed the spectrum of supernova neutrino events in ultrapure scintillators for a supernova exploding at 10 kpc from the Earth. We have examined the capability to distinguish the various detection channels and we have quantified the uncertainties in this type of detectors.  As discussed in the introduction, a major reason of specific interest  for a future supernova   is the possibility to observe neutral current interactions of neutrinos. We have investigated the three possible reactions of detection in ultrapure scintillators, namely: 1)~the elastic scattering with protons, 2)~the 15.11 MeV $\\gamma$ de-excitation line, 3)~the proton knockout channel. Our conclusions are as follows:   The first reaction is characterized by the larger number of expected events in all the detectors; however  the number of detectable events is strongly limited by the energy thresholds. The uncertainty on the total number of elastic scattering on protons, due to the proton structure, amounts to the 20\\%; moreover in the same energy region where this reaction can be observed there are also the indistinguishable events due to the elastic scattering with electrons and those due to NC and CC proton knockout. In other words, all we can observe is the total number of events collected in the energy region from the detector threshold to the threshold of the IBD signal, about $1.8$ MeV. In this detection window, % we have found that a fraction of 8\\% (Borexino), the 7\\% (KamLAND), the 4\\% (SNO+) of the signal is due to the other channels and this uncertainty is small but irreducible. We have also seen that, in the case that  the the energy of non-electronic neutrinos is low, the number of events due to this reaction is too small to permit the investigation of the $\\nu_x$ spectrum at the level discussed in~\\cite{dasg}.  The gamma line due to neutrino-induced $^{12}$C de-excitation is in principle easier, giving a signal at a high energies; its detection does not require the extreme performances at very low energies are not needed. However,  in the same region of the spectrum where this line is visible we will have also  positrons due to IBD reaction; thus, the efficiency to tag the concomitant neutron will be of crucial importance to   identify cleanly a sample of this NC reaction. While this concern is not a severe issue for the type of detectors we have considered in this work, it is much more relevant for future scintillators with a much larger mass and limited performances at low energies.  Finally, we have shown that the proton knockout will give a comparably small number of NC events."
    },
    "1402/1402.6212_arXiv.txt": {
        "abstract": "This is the second in a series of papers studying the astrophysics and cosmology of massive, dynamically relaxed galaxy clusters. The data set employed here consists of \\Chandra{} observations of 40 such clusters, identified in a comprehensive search of the \\Chandra{} archive for hot ($kT\\gtsim 5\\keV$), massive, morphologically relaxed systems, as well as high-quality weak gravitational lensing data for a subset of these clusters. Here we present cosmological constraints from measurements of the gas mass fraction, \\fgas{}, for this cluster sample. By incorporating a robust gravitational lensing calibration of the X-ray mass estimates, and restricting our measurements to the most self-similar and accurately measured regions of clusters, we significantly reduce systematic uncertainties compared to previous work. Our data for the first time constrain the intrinsic scatter in \\fgas{}, $7.4\\pm2.3$ per cent in a spherical shell at radii 0.8--1.2\\,$r_{2500}$ ($\\sim1/4$ of the virial radius), consistent with the expected level of variation in gas depletion and non-thermal pressure for relaxed clusters. From the lowest-redshift data in our sample, five clusters at $z<0.16$, we obtain a constraint on a combination of the Hubble parameter and cosmic baryon fraction, $h^{3/2}\\,\\Omegab/\\Omegam = 0.089\\pm0.012$, that is insensitive to the nature of dark energy. Combining this with standard priors on $h$ and $\\Omegab h^2$ provides a tight constraint on the cosmic matter density, $\\Omegam=0.27\\pm0.04$, which is similarly insensitive to dark energy. Using the entire cluster sample, extending to $z>1$, we obtain consistent results for \\Omegam{} and interesting constraints on dark energy: $\\Omegal=0.65^{+0.17}_{-0.22}$ for non-flat \\LCDM{} (cosmological constant) models, and $w=-0.98\\pm0.26$ for flat models with a constant dark energy equation of state. Our results are both competitive and consistent with those from recent cosmic microwave background, type Ia supernova and baryon acoustic oscillation data. We present constraints on more complex models of evolving dark energy from the combination of \\fgas{} data with these external data sets, and comment on the possibilities for improved \\fgas{} constraints using current and next-generation X-ray observatories and lensing data. ",
        "introduction": "\\label{sec:cosintro}  The matter budget of massive clusters of galaxies, and specifically the ratio of gas mass to total mass, provides a powerful probe of cosmology (\\citealt{White1993Natur.366..429W, Sasaki9611033, Pen9610090, Allen0205007, Allen0405340, Allen0706.0033}; \\citealt*{Allen1103.4829}, and references therein). In these systems, the mass of hot, X-ray emitting gas far exceeds that in colder gas and stars (e.g.\\ \\citealt{Lin0408557}; \\citealt*{Gonzalez0705.1726}; \\citealt{Giodini0904.0448, Dai0911.2230, Leauthaud1104.0928, Behroozi1207.6105}), and the gas mass fraction, \\fgas{}, is expected to approximately match the cosmic baryon fraction, $\\Omegab/\\Omegam$ (\\citealt{Borgani0906.4370}, and references therein). Hydrodynamic simulations of cluster formation indicate that the gas mass fraction at intermediate to large cluster radii should have a small cluster-to-cluster scatter and evolve little or not at all with redshift (\\citealt{Eke9708070, Kay0407058, Crain0610602}; \\citealt*{Nagai0609247}; \\citealt{Young1007.0887, Battaglia1209.4082, Planelles1209.5058}). Increasingly, as simulations have incorporated more accurate models of baryonic physics in clusters, in particular modeling the effects of feedback from active galactic nuclei (AGN) in cluster cores (e.g.\\ \\citealt{McNamara0709.2152}), they have become able to more reliably predict the baryonic depletion of clusters relative to the Universe as a whole. Combining such predictions with measurements of cluster \\fgas{} and constraints on $\\Omegab$, for example from cosmic microwave background (CMB) or Big Bang Nucleosynthesis (BBN) data and direct estimates of the Hubble parameter, provides a uniquely simple and robust method to constrain the cosmic matter density, $\\Omegam$. The pioneering work of \\citet{White1993Natur.366..429W} was among the first to show a clear preference for a low-density universe with $\\Omegam\\sim0.3$, a result which cluster \\fgas{} data continue to support with ever greater precision (\\citealt{Allen0205007, Allen0405340, Allen0706.0033, Ettori0211335, Ettori0904.2740}; \\citealt*{Rapetti0409574}), and which has been corroborated by a variety of independent cosmological data (e.g.\\ \\citealt{Percival0608635, Percival0907.1660, Spergel0603449,  Kowalski0804.4142, Mantz0709.4294, Mantz0909.3098, Dunkley0803.0586, Vikhlinin0812.2720, Rozo0902.3702, Blake1108.2635, Komatsu1001.4538, Hinshaw1212.5226, Suzuki1105.3470, Anderson1303.4666, Benson1112.5435, Hasselfield1301.0816, Planck1303.5076}).  Given a bound on the evolution of \\fgas{} from theory or simulations, the apparent evolution in \\fgas{} values measured from X-ray data can also provide important constraints on the cosmic expansion history and dark energy \\citep{Sasaki9611033, Pen9610090}. This sensitivity follows from the fact that derived \\fgas{} values depend on a combination of luminosity and angular diameter distances to the observed clusters, analogously to the way that type Ia supernova probes of cosmology exploit the distance dependence of the luminosity inferred from an observed flux. \\citet{Allen0405340} provided the first detection  of the acceleration of the cosmic expansion from \\fgas{} data, and more recently expanded and improved their analysis (\\citealt{ Allen0706.0033}, hereafter \\arsemf{}; see also \\citealt{LaRoque0604039, Ettori0904.2740}).  A key requirement for this work is that systematic biases and unnecessary scatter in the \\fgas{} measurements be avoided. This can be achieved by limiting the analysis to the most massive, dynamically relaxed clusters available. The restriction to relaxed systems minimizes systematic biases due to departures from hydrostatic equilibrium and substructure, as well as scatter due to these effects, asphericity, and projection \\citep{Rasia0602434, Nagai0703661, Battaglia1209.4082}. Similarly, using the most massive clusters minimizes residual systematic uncertainties associated with details of the hydrodynamic simulations, and simplifies the analysis by restricting it to those clusters for which \\fgas{} is expected to have the smallest variation with mass or redshift, and the smallest intrinsic scatter \\citep{Eke9708070, Kay0407058, Crain0610602, Nagai0609247,  Stanek0910.1599, Young1007.0887, Borgani0906.4370, Battaglia1209.4082, Planelles1209.5058, Sembolini1207.4438}. Moreover, the most massive clusters at a given redshift will be the brightest at X-ray wavelengths and require the shortest observing times.  This paper is the second of a series in which we study the astrophysics and cosmology of the most massive, relaxed galaxy clusters. The first installment (Mantz et~al.\\ 2014, in preparation, hereafter \\morphpaper{}) presents a procedure for identifying relaxed clusters from X-ray data based on their morphological characteristics, and identifies a suitable sample from a comprehensive search of archival \\Chandra{} data. In future work (\\profilespaper), we will investigate the astrophysical implications of our analysis of these clusters. This paper presents the cosmological constraints that follow from measurements of \\fgas{} for the cluster sample.  Our work builds directly on that of \\citet{Allen0205007, Allen0405340, Allen0706.0033}. Among our methodological improvements, three stand out as particularly important. First, the selection of target clusters has been automated (\\morphpaper{}), enabling straightforward application to large samples. Second, the cosmological analysis uses gas mass fractions measured in spherical shells at radii near $r_{2500}$,\\footnote{Defined as the radius within which the mean cluster density is 2500 times the critical density of the Universe at the cluster's redshift.} rather than \\fgas{} integrated at all radii $<r_{2500}$. The exclusion of cluster centers from this measurement significantly reduces the corresponding theoretical uncertainty in  gas depletion from hydrodynamic simulations.\\footnote{Improvements in the simulated physics, particularly the inclusion of feedback processes, have also been important in reducing this uncertainty (e.g.\\ \\citealt{Battaglia1209.4082, Planelles1209.5058}).} Third, the availability of robust mass estimates for the target clusters from weak gravitational lensing \\citep[][hereafter collectively Weighing the Giants, or \\wtg{}]{von-der-Linden1208.0597, Kelly1208.0602, Applegate1208.0605} allows us to directly calibrate any bias in the mass measurements from X-ray data, for example due to departures from hydrostatic equilibrium (e.g.\\ \\citealt{Rasia0602434, Nagai0609247, Battaglia1209.4082}) or instrument calibration (Applegate et~al., in preparation). In addition, our procedure employs blind analysis techniques (deliberate safeguards against observer bias) including hiding measured gas mass and total mass values until all analysis of the individual clusters was complete.  \\secref~\\ref{sec:cosmodata} reviews the selection of our cluster sample and basic X-ray data reduction (more fully described in \\morphpaper{}), as well as the additional analysis steps required to derive \\fgas{}. The resulting \\fgas{} measurements are presented in \\secref~\\ref{sec:fgresults}. The cosmology and cluster models we fit to the data are described in \\secref~\\ref{sec:model}, and \\secref{}~\\ref{sec:cosmores} presents the cosmological results. \\secref~\\ref{sec:changes} summarizes the differences between our work and \\arsemf{} (also discussed throughout \\secref{}s~\\ref{sec:cosmodata} and \\ref{sec:model}) and compares their cosmological constraints. In \\secref~\\ref{sec:future}, we discuss the potential for further improvements in \\fgas{} constraints from future observing programs targeting clusters discovered in upcoming surveys. We conclude in \\secref~\\ref{sec:cosmosummary}.  For the cosmology-dependent quantities presented in figures and tables, we adopt a reference flat \\LCDM{} model with Hubble parameter $h=H_0/100\\km\\second^{-1}\\Mpc^{-1}=0.7$, matter density with respect to the critical density $\\Omegam=0.3$, and dark energy (cosmological constant) density $\\Omegal=0.7$. However, our cosmological constraints are independent of the particular choice of reference (\\arsemf{} and \\secref~\\ref{sec:model}). ",
        "conclusions": ""
    },
    "1402/1402.4217_arXiv.txt": {
        "abstract": "We have made mass maps of three strong-lensing clusters, Abell 3827, Abell 2218 and Abell 1689, in order to test for mass-light offsets. The technique used is GRALE, which enables lens reconstruction with minimal assumptions, and specifically with no information about the cluster light being given.  In the first two of these clusters, we find local mass peaks in the central regions that are displaced from the nearby galaxies by a few to several kpc.  These offsets {\\em could\\/} be due to line of sight structure unrelated to the clusters, but that is very unlikely, given the typical levels of chance line-of-sight coincidences in $\\Lambda CDM$ simulations --- for Abell 3827 and Abell 2218 the offsets appear to be intrinsic.  In the case of Abell 1689, we see no significant offsets in the central region, but we do detect a possible line of sight structure: it appears only when sources at $z\\ga 3$ are used for reconstructing the mass.  We discuss possible origins of the mass-galaxy offsets in Abell 3827 and Abell 2218: these include pure gravitational effects like dynamical friction, but also non-standard mechanisms like self-interacting dark-matter. ",
        "introduction": "Our current understanding of the universe and its dynamics indicates that its major components are dark: cold dark-matter (CDM) and the so-called ``dark-energy''.  Unlike baryons, dark-matter interacts only gravitationally and provides the deep potential wells which are followed by the baryons. The baryons form clumps at these potential wells and cool down to form stars.  The standard $\\Lambda$CDM model explains a range of observed processes pretty well, from the angular power spectrum of the cosmic microwave background \\citep{2013arXiv1303.5075P} to the baryonic acoustic oscillations \\citep{2013MNRAS.433.1202S} in the large scale structure and the number counts of clusters.  However, the intrinsic properties and behaviour of dark-matter and dark-energy remain an open problem in cosmology.  In the picture of hierarchical structure formation in $\\Lambda$CDM model, galaxy-clusters are the most recently formed structures that are gravitationally bound. They are cosmic laboratories to test the laws of gravity, structure formations and the interaction of different species of particles. A galaxy cluster contains lots of galaxies --- tens to thousands, hot intra-cluster plasma visible in X-rays, a variety of of relativistic particles and finally dark-matter which dominates its mass budget.  Measuring the mass of the galaxy-cluster is an essential aspect of using the cluster to study many other things. There are several physical processes that enable one to measure the mass: the kinematics of cluster galaxies \\citep{2013ApJ...772...47S}, the hydrodynamics of hot gas emitting X-rays \\citep{2009ApJ...692.1033V}, and gravitational lensing.  Lensing is particular interesting, because it relies only on gravity and does not itself require any luminous objects in the cluster being studied. One of the questions that lensing can address is how well the luminous matter traces the distribution of total mass. Deviations, or lack thereof, from the mass-follows-light hypothesis will provide important information about the physical processes going in within clusters. The first lensing-based detection of deviations from mass-follows-light goes back to the late 1990's \\citep{1998AJ....116.1541A} but the observation that generated a wide interest in these deviations was that of the Bullet Cluster \\citep{2006ApJ...648L.109C}, which showed unambiguously that dark matter is quite collisionless compared to the gas phase baryonic matter \\citep{2008ApJ...679.1173R}. While the properties of dark matter are probably not the only reason for deviations from mass-follows-light in galaxy clusters, dark matter self-interaction cross-section and how to optimally extract it from observations is an exciting avenue of research \\citep{2013arXiv1310.1731H,2013MNRAS.433.1517H}.  This work uses strong gravitational lensing to look for deviations from mass-follows-light, i.e. it explores the correspondence on the sky between the dark-matter peaks with the galaxies in the central parts of three galaxy clusters, Abell 3827, 2218 and 1689. These clusters are very different from each other in morphology and redshift. As we discuss in Section~\\ref{disc}, some deviations we find may be due to the non-standard properties of dark matter, but others could be the result of superimposed substructure, or hydrodynamics within the cluster.  We use GRALE \\citep{2006MNRAS.367.1209L,2007MNRAS.380.1729L}, a strong-gravitational lensing tool to reconstruct the mass map of the clusters.  There is no overall parametric form for the mass distribution, but rather an adaptive grid.  Other than the redshift, no information about the cluster is required as input, not even its location or morphology.  This makes GRALE well-suited to reconstruction of mass maps before comparison with light. ",
        "conclusions": "\\label{disc}  Gravitational lensing offers a unique opportunity to study the distribution of matter in clusters of galaxies. Free-form reconstruction methods take full advantage of this.  Our synthetic tests show that GRALE recovers the mass distribution well, and the concomitant errors provide a reliable guide to assessing the significance of various mass features. The test case in Figure \\ref{fig:plummerout} and \\ref{fig:plummer5out} shows no spurious offsets in the mass maps.  Reconstructions of the three real lensing clusters indicate some curious features.  In two clusters we see offsets between the optical light and the nearest mass concentrations.  The form of the offsets is not resolved: they could be distinct peaks in the projected mass distribution; or they could be spurs that extend from a peak that itself coincides with the galaxy light; or the offsets could very lopsided dark halos around galaxies.  (We emphasize that not all offsets seen in the reconstructed mass maps are significant, but only those that pass the statistical significance tests.)  A caveat to bear in mind is the assumption that the observed image positions are accurate. Because lensed images are often faint, have low surface brightness and are superimposed on brighter cluster galaxies, image identification is not always straightforward. It is thus conceivable that some images have been misidentified.  But assuming the image identifications are all valid, confirmation by independent techniques is desirable.  Lens reconstruction methods not assuming light traces mass in some way include Lensview \\citep{2006MNRAS.372.1187W}, LensPerfect \\citep{2008ApJ...681..814C} and PBL \\citep{2008ApJ...687...39D} and any of these would be suitable.  If the mass/galaxy offsets are confirmed, they would lead to interesting conclusions about the nature of clusters and dark-matter.  In general, several reasons for offsets are possible. Superimposed, but dynamically unrelated line of sight structures could contribute lensing mass, with no apparent associated light, especially if the structures are considerably further away from us than the main lensing cluster. However, we argue that the offset in A3827,  is not due to the line of sight structure because of the very low redshifts of the sources. In A2218 line of sight structures are also unlikely to be the cause because only a very concentrated and massive los structure can contribute significantly in the vicinity of a massive clump within a cluster. Such chance superposition are expected to be rare.  Line of sight structures are more likely to make a contribution away from mass concentrations within the cluster, where cluster projected densities are lower. This can be illustrated with dark-matter N-body simulations. The blue lines in Figure~\\ref{fig:simulation} are the isodensity contours of the total projected mass in a cylinder centered on a halo whose virial radius is the radius of the window, while the red lines are the contours of the projected mass inside the virial sphere of the cluster. We caution that  these plots were made with a limited line of sight depth of about comoving 90 Mpc (Simulations courtesy J\\\"urg Diemand; \\cite{2004MNRAS.352..535D}). The black contours mark regions where the fractional mass excess due to the line of sight structures (and not the mass within the virial sphere) amount to $25\\%$ of total.  The top two panels show examples where the contribution from the los material is typical, while the bottom two panels present two cases with the most contribution (out of a total of 100 lines of sight). Even though the length of the cylinder is not large, the plots show that los structures cannot make a significant contribution where the cluster density is high. However, such structures can make a significant contribution at some distance away from the cluster centre.  In A1689 we might be seeing such a line of sight structure. After subtracting the mass reconstruction based on high-$z$ sources (HRS) from that based on low-$z$ sources (LRS) we see a mass concentration about 30 arcsec, or 100 kpc from cluster center. It is statistically significant (it contributes to the tail of the distribution shown in Fig.~\\ref{fig:a1689_los} which extend beyond the right edge of the plot) but is not associated with bright cluster galaxies. We interpret it as arising from  a structures between the $z\\approx 2$ and $3$.  If not line of sight structure, what else can be responsible for mass-light offsets seen in A3827 and A2218?  Offsets could be intrinsic to the cluster, and be due to manifestations of known physics, like gravity, and hydrodynamics of the gas, or new physics, such as self-scattering of dark matter. Offsets in merging clusters have been observed, but mostly between the dark matter and the X-ray emitting gas components \\citep{2006ApJ...648L.109C,2013MNRAS.429..833H,2012ApJ...758..128C}. In the outskirts of Abell 2744 a separation between dark matter and galaxy components is also seen \\citep{2011MNRAS.417..333M}, and in the merging cluster CL0152-1357 an offset between Sunyaev-Zel'dovich effect and X-ray peaks has been detected \\citep{2012ApJ...748...45M}.  Most of these offsets are on larger scales then what we detect in this work. For smaller scale offsets early stage mergers are probably not the cause, and different set of causes has to be considered.  One of the possibly relevant gravitational effects is the oscillation or wobbling of a galaxy, such as a BCG around the bottom of the gravitational potential. This has been observed in a sample of galaxy clusters as a displacement of the BCG from the lensing centroid \\citep{2012MNRAS.426.2944Z}. The distribution is displacements is wide, and peaks at roughly 10~kpc.  Whether this is a likely explanation for the offsets in A3827 and A2218 is yet to be determined---the observed offsets are not for central cluster galaxies.  It is less likely, but still possible that the offsets are a consequence of tidal effects. These would strip the material from the galaxy symmetrically in the leading and trailing directions. Since the offsets in A3827 and A2218 do not show such symmetry, tidal effects are probably not the main cause.  Dynamical friction would create an asymmetric structure and would preferentially distort the distribution of dark matter and not stars if the former has a more extended distribution. A numerical simulation would be required to test this possibility.  The formation of a galaxy cluster is a complex process involving hydrodynamics of gas. It is possible that  star formation induced by galaxy mergers within clusters would result in stars and dark matter halos offsets.  Finally, if dark-matter has non-negligible self-interaction cross-section, dark-matter particles of the galaxy halo would experience a drag force as the galaxy moves within the halo of the cluster. The nature of the resulting dark-matter features induced by these interactions may be consistent with those observed in A3827 and A2218, but detailed simulations are required \\citep{2014MNRAS.437.2865K}."
    },
    "1402/1402.3601_arXiv.txt": {
        "abstract": "Kinetic inductance detectors (KIDs) show promise as a competitive technology for astronomical observations over a wide range of wavelengths.  We are interested in comparing the fundamental limitations to the sensitivity of KIDs with that of transition edge sensors (TESs) at millimeter wavelengths, specifically over the wavelengths required for studies of the Cosmic Microwave Background (CMB). We calculate the total fundamental noise arising from optical and thermal excitations in TESs and KIDs for a variety of bath temperatures and optical loading scenarios for applications at millimeter wavelengths. Special consideration is given to the case of ground-based observations of 100 GHz radiation with a 100 mK bath temperature, conditions consistent with the planned second module of the QUBIC telescope, a CMB instrument \\cite{Battistelli2011}. Under these conditions, a titanium nitride KID with optimized critical temperature pays a few percent noise penalty compared to a typical optimized TES. ",
        "introduction": "Two types of superconducting, incoherent detectors of millimeter and sub-millimeter radiation have undergone rapid development in the past decade: transition edge sensors and kinetic inductance detectors. TES arrays of 1000s of pixels with micromachined thermal isolation structures have been developed which approach background-limited sensitivity for both ground-based and space-based observations of faint sources such as the CMB.  KID arrays with similar pixel counts have been deployed in instruments such as MUSIC, but they have not yet reached these same sensitivity levels.  However, the simplicity of fabricating and multiplexing KIDs has sparked considerable interest and naturally invites comparison with TES arrays in terms of their ultimate sensitivity.  For an optimized TES, the fundamental noise limits arise from two sources: photon noise from the astrophysical sources or backgrounds under observation, and thermal noise caused by fluctuations in thermal carriers (typically phonons) passing through the weak thermal isolation link from the bolometer's absorbing structure to the thermal bath. For KIDs, photon noise is also present, as is generation-recombination (g-r) noise, caused by fluctuations in the number density of quasiparticles in the device from optical or thermal excitations. It is similar to g-r noise in a photoconductor, caused by fluctuations in the number of charge carriers in a semiconductor. Currently, KIDs are limited by other noise sources, such as two-level system noise in the dielectric substrates, but rapid progress is being made in understanding and reducing these noise contributions \\cite{Zmuidzinas2012}.  In this paper we compare the `fundamental noise' limits of KID and TES detectors, particularly as applied to observations of the CMB. We study their performance in both ground-based and space-based optical loading conditions and with bath temperatures below 250 mK, typical of the cryogenic systems used for CMB applications.  We begin by describing the details of the dominant detector noise contributions for KIDs and TESs, as well as our optimization scheme and our assumptions for each detector type.  Finally, we present our comparison results. ",
        "conclusions": ""
    },
    "1402/1402.4810_arXiv.txt": {
        "abstract": "{ The molecular phase of the interstellar medium (ISM) in galaxies offers fundamental insight for understanding star-formation processes and how stellar feedback affects the nuclear activity of certain galaxies. We  present here Large Millimeter Telescope spectra obtained with the Redshift Search Receiver, a spectrograph that covers simultaneously the 3 mm band from 74 to 111 GHz with a spectral resolution of around 100 km\\,s$^{-1}$. Our selected galaxies, have been detected previously in HCN, and have different degrees of nuclear activity --- one normal galaxy (NGC~6946), the starburst prototype (M~82) and two ultraluminous infrared galaxies (ULIRGs, IRAS~17208-0014 and Mrk~231). We plotted our data in the HCO$^+$/HCN vs. HCN/$^{13}$CO diagnostic diagram finding that NGC~6946 and M~82 are located close to other normal galaxies; and that both IRAS~17208-0014 and Mrk~231 are close to the position of the well known ULIRG Arp~220 reported by~\\citet{Snell2011}. We found that in Mrk~231 -- a galaxy with a well known active galactic nucleus -- the HCO$^+$/HCN ratio is similar to the ratio observed in other normal galaxies. } ",
        "introduction": "The starburst-- active galactic nucleus (AGN) connection is one of the most interesting and evolving topics of  modern astronomy. Learning about nuclear activity in galaxies implies segregating the different components present in the nuclei of galaxies, plus allowing for a range of ages, metal content and environmental differences. For example, when dealing with AGNs and nuclear starbursts (SBs)  one has to take into account that AGNs tend to appear preferentially in early type galaxies while SBs tend to appear in Hubble types later than Sbc. This implies that -- on average -- physical parameters like metallicity or the age of the underlying stellar population are different in AGNs and SBs. In general, Seyfert galaxies have metal content higher than the SB galaxies, thus the properties of the star formation in SB nuclei are not necessarily similar to the star formation in Seyfert nuclei. Part of the problem is that galactic nuclei, where these processes take place, can be very dusty. This compromises the interpretation of short-wavelength observations, and makes a multi-wavelength approach imperative. One possible solution is to take advantage of sub-millimeter and millimeter (mm) spectra, only recently possible thanks to new facilities. The mm-wavelength spectral surveys are beginning to provide sensitive measurements of an astounding number of molecular transitions in starbursts and AGNs \\citep[e.g.][]{Burillo2010,Aladro2011,Aladro2013}. Such transitions allow the detailed study of molecular chemistry in galactic nuclei, interpreted through theoretical models only now starting to achieve the necessary sophistication \\citep[e.g.][]{Papadopoulos2010,Kazandjian2012,Bayet2012}. The mm spectral line ratios provide diagnostics to interpret the relative contributions of different physical environments in and around molecular clouds in the nuclear regions of galaxies. Depending on the strengths of AGN accretion  and star formation, molecular chemistry can be driven by Photo-Dissociated Regions (PDRs), Cosmic-Ray Dominated Regions (CRDRs), X-ray Dominated Regions (XDRs), shock-dominated regions (or mechanically heated regions), and dense shielded regions which lie deep in the molecular cloud ensembles. One of the main parameters that drives the evolution of a given galaxy is the star formation rate (SFR) which is related to the presence of dense shielded cores with n(H$_2$)$>3\\times 10^4$ cm$^{-3}$. These regions are traced by molecules with high dipole moment such as HCN and CS. In fact the HCN luminosity in CO bright galaxies seems to correlate very well with the infrared luminosity, a common tracer of the star formation activity \\citep{Gao2004a,Gao2004b}. Lines in the 3mm spectral region such as CO, HCN, HNC, HCO$^+$, and CS have been used \\citep[e.g.][]{Kohno2001,Carpio2006,Snell2011} to reveal and separate  PDRs and CRDRs, associated with SF and intense starbursts, from XDRs, typically associated with AGNs \\citep[e.g.][]{Papadopoulos2010,Meijerink2011}. It is interesting to see that the old problem of determining the source of energy in the nucleus of nearby galaxies -- star formation or accretion to black holes -- is still present in recent studies of molecular chemistry. Rare isotopes such as C$^{18}$O and $^{13}$CO may trace rapid ISM enrichment and in some cases they have been related to the presence of WR stars or Type II SNe (e.g., Costagliola et al. 2011). ",
        "conclusions": "\\label{subsec:2} The observations discussed in this paper include M~82 which was observed with the RSR at the LMT during the first-light scientific demonstration observations phase in June 2011, and NGC~6946, IRAS~17208-0014 and Mrk~231 that where observed during the summer of 2013. In the observed 3 mm spectra of these galaxies we detected several lines including, in some cases, lines as weak as CN which traces very dense molecular gas (n(H$_2)\\sim 10^6$cm$^{-3}$). However, we are reporting only our results on the HCN,  HCO$^+$ and $^{13}$CO lines. In Fig.~\\ref{fig:TheFigure} we plotted our results in the HCO$^+$/HCN vs. HCN/$^{13}$CO diagnostic diagram. In that diagram the horizontal axis is an indicative of the dense gas fraction, where galaxies with high fraction of dense gas, N(H$_2)\\sim 10^5$ cm$^{-3}$, are located to the right, and the vertical one is related with the presence of an AGN where those galaxies with HCO$^+$/HCN much greater than 1 are better explained by models that include X-ray dominated regions ~\\citep[e.g.][]{Meijerink2011}. M~82 and NGC~6946 are close to other normal galaxies reported by~\\citet{Snell2011}. In that paper they also observed M~82 (marked with a D in the Figure) with the same receiver but mounted in the FCRAO 15 m antenna. The differences between  the HCN, $^{13}$CO ratio can be explained by the different beam sizes between the LMT and the FCRAO and the relative distribution of the dense gas traced by HCN and the more diffuse gas traced by $^{13}$CO. The two ultra--luminous galaxies (ULIRGs) observed,  IRAS~17208-0014 and Mrk~231, are close to the position of the well known ULIRG Arp 220, indicating the presence of high amounts of dense gas in both galaxies. Notice that Mrk~231, a well known active galaxy classified as a Seyfert~1, has not an HCO$^+$ enhancement as predicted by the theoretical models. By using all the available data of this galaxy, which cover all the electromagnetic spectrum, we expect to better understand this discrepancy.  We are witnessing the capability of the RSR--LMT system through the first results of an ongoing project where we will observe hundreds of galaxies with different degree of central activity -- from HII galaxies, starbursts, LINERS, Seyfert 1--2, and luminosities -- from normal galaxies to luminous and ultra luminous infrared galaxies, and we will be able to relate with high statistical significance the properties of nuclear activity with the complex molecular chemistry.  \\begin{acknowledgement} We thank the project head of LMT David Hughes and the INAOE director Alberto Carrami\\~nana for their continuous support to the project. We also appreciate the work carried out by the LMT staff.  \\end{acknowledgement}"
    },
    "1402/1402.7022_arXiv.txt": {
        "abstract": "The origin of gamma-ray burst (GRB) prompt emission, bursts of $\\gamma$-rays lasting from shorter than one second to thousands of seconds, remains not fully understood after more than 40 years of observations. The uncertainties lie in several open questions in the GRB physics, including jet composition, energy dissipation mechanism, particle acceleration mechanism, and radiation mechanism. Recent broad-band observations of prompt emission with Fermi sharpen the debates in these areas, which stimulated intense theoretical investigations invoking very different ideas. I will review these debates, and argue that the current data suggest the following picture: A quasi-thermal spectral component originating from the photosphere of the relativistic ejecta has been detected in some GRBs. Even though in some cases (e.g. GRB 090902B) this component dominates the spectrum, in most GRBs, this component either forms a sub-dominant ``shoulder'' spectral component in the low energy spectral regime of the more dominant ``Band'' component, or is not detectable at all. The main ``Band'' spectral component likely originates from the optically thin region due to synchrotron radiation. The diverse magnetization in the GRB central engine is likely the origin of the observed diverse prompt emission properties among bursts. ",
        "introduction": "Although discovered much earlier than the afterglow, the emission of gamma-ray bursts (GRBs) themselves, usually called the GRB {\\em prompt emission}, is still poorly understood. The main uncertainties lie in several open questions in the GRB physics \\cite{zhang11cr}: \\begin{itemize} \\item What is the composition of the GRB jet? \\item What is the energy dissipation mechanism in the jet? \\item Where is the region of energy dissipation and prompt emission? \\item What is the particle acceleration mechanism in the energy dissipation region? \\item What is the radiation mechanism of the particles? \\end{itemize} Among these, the answer to the first question is more fundamental, which decides or at least strongly affects the answers to the rest of the questions.  The traditional ``fireball'' model envisages a thermally-driven explosion, with gravitational energy released during a catastrophic event (such as massive star core collapse or merger of two compact stars) being deposited in the form of thermal heat at the base of the central engine. This fireball then expands under its own thermal pressure and gets accelerated to a relativistic speed, if baryon contamination in the ejecta is low enough \\cite{paczynski86,goodman86}. Most of the thermal energy is converted to the kinetic energy of the outflow \\cite{shemi90}, with some thermal energy released as photons at the photosphere \\cite{meszarosrees00}. The kinetic energy of the outflow is further dissipated in the internal shocks \\cite{rees94} and in the external (forward and reverse) shocks as the ejecta is decelerated by the ambient medium \\cite{rees92,meszarosrees93}. The former powers the GRB prompt emission, while the latter powers the afterglow \\cite{meszarosrees97,sari98}.  In contrast to the matter-dominated fireball picture, an opposite opinion is that the entire GRB phenomenology is electromagnetic. This is the electromagnetic (EM) model of GRBs \\cite{lyutikov03}. Here the magnetic field plays the dominant role in ejecta dynamics. Little baryons are loaded from the central engine, so that an extremely high value of $\\sigma$ (ratio between Poynting flux and matter flux, or ratio between comoving magnetic energy density and rest mass energy density) is invoked. It is conjectured that $\\sigma$ remains extremely high ($\\sim 10^5 - 10^6$) at the deceleration radius, so that the ``sub-Alfvenic'' condition $\\sigma > \\Gamma^2-1$ ($\\Gamma$ is the bulk Lorentz factor of the outflow) is satisfied. Within this scenario, the magnetic energy is directly converted to the particle energy and radiation energy through current instability around the deceleration radius. Such a scenario requires that the ejecta maintains high magnetization throughout, which is very difficult to achieve in reality.  More realistically, the GRB ejecta may carry both a matter component and a magnetic component. This results in a wide range of magnetohydrodynamic (MHD) models of GRBs \\cite{spruit01,drenkhahn02,giannios08,zhangyan11,mckinney12}. A real question is how high the $\\sigma$ parameter is. Since $\\sigma$ is a function of radius (magnetic acceleration and dissipation can constantly change $\\sigma$ in the flow), a more relevant question would be how high the $\\sigma$ parameter is at the GRB emission site. The complication is that one is not sure about the location of the GRB prompt emission radius $R_{\\rm GRB}$ (see next section). The values of $\\sigma(R_{\\rm GRB})$ and $R_{\\rm GRB}$ are mutually dependent (and also depend on a third parameter, i.e. $\\sigma_0$ at the central engine). Different GRB prompt emission models invoke different combinations of these parameters.  Observationally, the currently available information of GRB prompt emission includes the following aspects. A successful model should be able to interpret them. \\begin{itemize} \\item GRB lightcurves are extremely diverse \\cite{fishman95}. Some are spiky, while some others are smooth. Some have multiple (sometimes well separated) spiky episodes, while some others have one or two smooth pulses. Some are very long, while some others are very short. At least some lightcurves can be decomposed as the superposition of a fast and a slow component\\cite{gao12}. \\item GRB spectra can be well characterized. For most GRBs, both the time-integrated and time-resolved spectra can be described by a so-called ``Band'' function, i.e. a smoothly-joint broken power law function characterized by three parameters: the low-energy and high-energy photon indices $\\alpha$ and $\\beta$, as well as the peak energy $E_p$ in the energy spectrum \\cite{band93}. Recent Fermi LAT/GBM observations suggest that such a component dominates in nearly 7 orders of magnitude in energy in most GRBs \\cite{abdo09a,zhang11}. On the other hand, growing evidence suggests that this phenomenological function cannot fully account for the observations, and two additional spectral components, i.e. a quasi-thermal bump and a high-energy power-law component, are needed to fit the time-resolved data of bright GRBs \\cite{zhang11,ryde10,guiriec11,axelsson12,ackermann11}. \\item The peak energy $E_p$ varies in a wide range, from as high as 15 MeV in GRB 110721A \\cite{axelsson12} to below 10 keV in GRB 060218 \\cite{campana06}. It can evolve rapidly within a same burst, and two patterns can be identified: hard-to-soft evolution throughout a pulse or $E_p$ intensity tracking \\cite{lu10,lu12}. \\item GRB polarization measurements have been made for some GRBs. Even though not with very high significance, evidence of high linear polarization degree for at least some bright GRBs has been collected \\cite{yonetoku11,yonetoku12}. \\end{itemize} The available data do not have ``smoking gun'' criteria to differentiate among the different models. Nonetheless, some available data already brought useful clues, serving as ``finger prints'' to diagnose the underlying physics of GRB prompt emission. ",
        "conclusions": ""
    },
    "1402/1402.2257.txt": {
        "abstract": "Prior to recombination, Silk damping causes the dissipation of energy from acoustic waves into the monopole of the Cosmic Microwave Background (CMB), resulting in spectral distortions. These can be used to probe the primordial scalar power spectrum on smaller scales than it is possible with CMB anisotropies.  An enhancement of power on these scales is nevertheless required for the resulting distortions to be detectable by future experiments like PIXIE.  In this paper, we examine all 49 single-field inflation models listed by Martin et al. in the \\textit{Encyclopaedia Inflationaris}~\\cite{Martin:2013tda} and find that only one of these may lead to a detectable level of distortions in a tuned region of its parameter space, namely the original hybrid model.   Three effective multi-field scenarios are also studied: with softly and suddenly turning trajectories, and with a mild waterfall trajectory. Softly turning trajectories do not induce  distortions at any detectable level, whereas a sudden turn in the field space or a mild waterfall trajectory predicts a peak (plus damped oscillations in the sudden turn case) in the scalar power spectrum, which can lead to an observable amount of CMB distortions.  Finally, another scenario leading to potentially detectable distortions involves a curvaton whose blue spectrum is subdominant on CMB angular scales and overtakes the inflaton spectrum on smaller scales. In this case however, we show that the bounds from ultra compact minihaloes are not satisfied.  Expectations for an ultimate PRISM-class experiment characterized by an improvement in sensitivity by a factor of ten are discussed for some models. ",
        "introduction": "During the last two decades, observations of the Cosmic Microwave Background (CMB) anisotropies have been allowing for an increasingly accurate determination of the energy content and the shape of the Universe, as well as of the primordial density perturbations, that are at the origin of large scale structure formation~\\cite{Ade:2013lta}.  In order to match observations, these primordial inhomogeneities can be described as Gaussian and have a nearly scale-invariant power spectrum, which is a generic prediction of inflation models in which the Universe undergoes an early phase of exponentially accelerated expansion.   The most recent experiments like the Atacama Cosmology Telescope~\\cite{Das:2013zf}, and the South Pole Telescope~\\cite{Story:2012wx} have measured the CMB angular power spectrum up to multipoles $l=l_{\\rm max} \\sim 3000$, while Planck~\\citep{Planck:2013kta} observes at the same time large scales up to the lowest multipoles as well.  This range of angular scales probes wavelength modes of primordial density perturbations within the window $ k_{l \\rm max}=2 \\times 10^{-4} \\Mpc^{-1} \\lesssim k \\lesssim 0.2 \\Mpc^{-1}=k_{l \\rm min}$.  In the context of inflation, the largest have exited the Hubble radius about 60-efolds before the end of inflation and have since been frozen until they recently have re-entered the horizon.  CMB anisotropy experiments can therefore probe about $7\\approx\\ln  k_{\\ell \\rm max}/k_{\\ell \\rm min}$ e-folds of inflationary expansion.  Earlier and later periods of the inflationary epoch remain inaccessible to observations, respectively either because the corresponding scales are still super-horizon today, or because they have led to acoustic oscillations that are highly suppressed due to Silk damping.   Recently the BICEP2~\\cite{Ade:2014xna} experiment has claimed the discovery of B-mode polarization in the CMB spectrum induced by primordial gravitational waves corresponding to a tensor to scalar ratio $r = 0.20 ^{+0.07}_{-0.05} $.  These results are nevertheless still controversial (see e.g.~\\cite{Mortonson:2014bja}) and thus we adopt a conservative approach and do not consider BICEP2 constraints on the studied models in this paper.  Other signals and observational techniques must therefore be envisaged in order to extend the available information on the inflationary epoch by several e-folds. One example are the 21cm angular power spectra from the dark ages and the subsequent reionisation (see e.g Refs.~\\cite{Pritchard:2011xb,Pritchard:2010pa,Furlanetto:2009qk,Clesse:2012th}).  Measuring distortions of the CMB black-body radiation is another, less futuristic, interesting possibility.  Spectral distortions can be the consequence of energy injection into the primeval plasma through several processes~\\cite{Sunyaev:2013aoa,Chluba:2013wsa,Chluba:2011hw}: the decay or the annihilation of relic particles, the evaporation of primordial black holes, the interaction of CMB photons with non-relativistic electrons and baryons after recombination~\\cite{Chluba:2011hw}, as well as recombination lines, reionisation and structure formation. Here, we are mainly interested in the energy stored in sound waves and that is dissipated due to Silk damping into the monopole, which creates specific distortions of the CMB black-body spectrum~\\cite{Chluba:2013dna,Khatri:2012rt,Chluba:2012we,Chluba:2012gq,Khatri:2013dha}.  Measuring such distortions can therefore in principle provide some information on the primordial power spectrum of density perturbations, on scales much smaller than the ones probed by CMB anisotropy experiments.  Future dedicated CMB distortion experiments could probe modes in the range $8 \\lesssim k \\lesssim 10^4 \\Mpc^{-1}$, potentially extending the period of inflation that is accessible to observations up to 17 e-folds~\\cite{Khatri:2013xwa}.  Any perturbations of the thermodynamic equilibrium between photons and baryons during the tightly coupled regime and towards its end can induce CMB spectral distortions.  At very early times (redshifts $z\\gtrsim 2 \\times 10^6$), thermalization is very efficient and erases distortions caused by any amount of energy injection.  However, at lower redshifts, spectral distortions may survive and still be observable today.  It is common practice to distinguish between distortions characterised by a frequency-dependent chemical potential $\\mu(\\nu)$, the so-called $\\mu$-distortions, %and distortions characterised respectively by a reduction/increase of the intensity at low/high frequencies, the so-called $y$-distortions. and distortions of the $y$-type mostly associated to SZ clusters.  For completeness, one one should also consider partially Comptonized intermediate $i$-distortions, as discussed in Refs.~\\cite{Khatri:2012tw,Khatri:2013dha}. The current 95$\\%$C.L. limits on $\\mu$- and $y$-distortions result from the COBE-FIRAS experiment~\\cite{Fixsen:1996nj} and are $y<1.5 \\times 10^{-5}$ and $\\mu < 9 \\times 10^{-5}$.  Recently, the Primordial Inflation Explorer (PIXIE) \\cite{Kogut:2011xw} experiment has been proposed. It is one of its objectives to improve these limits by about three orders of magnitude.    Further in the future, the next to next generation of experiments, of the class of the Polarized Radiation Imaging and Spectroscopy Mission (PRISM) \\cite{Andre:2013nfa,Andre:2013afa}, will have the ability to improve the sensitivity to spectral distortions by about one order of magnitude compared to PIXIE.  Some work has been done recently to forecast possible constraints on the amplitude of the power spectrum of curvature perturbations, its spectral index and the running of the latter, for various experimental configurations~\\cite{Chluba:2012we,Khatri:2013dha,Chluba:2013dna}.  As a result, it is not expected that the PIXIE experiment, by probing only spectral distortions, will reach the sensitivity to improve current constraints on the spectral index in models for which its value on usual CMB scales can be extended to scales probed by CMB distortions.  But it is still unclear whether new constraints will be established on specific inflation models that predict an increase of power on scales smaller than the ones probed by CMB anisotropy experiments. It has also not yet been studied whether a curvaton, which is an alternative to the inflaton as the field responsible for the generation of the curvature power spectrum, can lead to observable CMB distortions.  In this paper, we adopt a model-oriented approach.  We have studied all the single-field inflation models listed in Ref.~\\cite{Martin:2013tda} and have established a short list of models predicting an increase of power on small scales for some part of their parameter space. For each of these we have studied whether they can lead to an observable level of CMB distortions or not, and we have determined the regions in the parameter space where this occurs.  For this purpose, we have modified the \\idistort template~\\cite{Khatri:2012tw},  as well as the \\Greens package~\\cite{Chluba:2013vsa} for Markov-Chain-Monte-Carlo (MCMC) parameter estimation, in order to include any shape of the power spectrum of curvature perturbations, and for all the relevant models we have calculated the spectrum of CMB distortions.  In addition to single field models, we have considered three types of multi-field scenarios: i) an effective model where the field trajectories perform a turn with a constant angular velocity at some point of the inflationary evolution, ii) an effective model where trajectories perform a sudden turn in the field space, and iii) an effective hybrid model for which between $30$ and $60$ e-folds of expansion are realised during the classical field evolution in the final waterfall phase.  Finally, we study a curvaton scenario that can possibly induce a detectable level of distortions.  We explore the signatures of each model on the CMB black-body spectrum and determine whether they could be measured by PIXIE.  For some models we also discuss expectations for a PRISM-class experiment.  The paper is organized as follows: In Sec.~\\ref{sec:distintro}, the different types of distortions are introduced.  Sec.~\\ref{sec:sensitivity} is dedicated to the expected sensitivity of the PIXIE experiment  and to forecasts on distinguishing different shapes of the scalar power spectrum.  In Sec.~\\ref{sec:single_field}, single field models that possibly lead to CMB distortions observable by PIXIE are studied.   We investigate three effective scenarios of multi-field inflation in Sec.~\\ref{sec:multi_field} and focus on a specific curvaton model in Sec.~\\ref{sec:curvaton}. To exemplify the benefit of improving the spectrometer's sensitivity beyond the level proposed for PIXIE, in Sec.~\\ref{sec:PRISM}, we investigate how a PRISM-class probe could be used to determine the running of the spectral index or a possible sub-dominant contamination of the scalar power spectrum from a curvaton. Finally, we summarize our findings and discuss the possible implications for the dedicated CMB distortion experiments of the next generation in the Conclusions (Sec.~\\ref{sec:ccl}). ",
        "conclusions": "\\label{sec:ccl}  We have performed the first exhaustive model-oriented analysis of the observational prospects for inflation in view of future observations of spectral distortions.  These are produced prior to recombination, when acoustic waves of Silk-damping scales dissipate their energy into the temperature monopole.   We find that only very few models can lead to distortions at a level that is detectable by the future PIXIE experiment.  After having introduced a methodology for the model selection, all 49 single-field inflation models listed in the \\textit{Encyclopaedia Inflationaris} by Martin et al. in Ref.~\\cite{Martin:2013tda}  have been considered.  We find that only one of them can satisfy simultaneously the Planck constraints on the primordial scalar power spectrum and generate a significant increase of power on smaller distortion scales, leading to a detectable level of intermediate and $\\mu$-type distortions, while the $y$-type component should be dominated by other sources, such as the thermal SZ effect. This one scenario of interest is the original hybrid model with an instantaneous waterfall phase but super-Planckian field values.   The increase of power on small scales results from the field dynamics evolving from the field dominated towards the false vacuum dominated inflationary regime, where the spectrum is generically blue.   Finally, it is important to note that the region of parameter space  that is interesting in the present context appears strongly tuned. Therefore, it would be rather fortuitous, if distortions from single-field inflation were detected by PIXIE.   The situation could improve in the future in the context of a PRISM-class experiment, characterized by a ten times better sensitivity.  In this case, not only models with a sufficient increase of power on distortion scales could be tested but also models predicting a large running or, more generally, models with a substantial change of the amplitude of the scalar power spectrum between distortion and CMB anisotropy scales.  For our model selection, we find that some parameter ranges of the Non-Canonical K\\\"ahler and the Running Mass inflation models can be tested by a PRISM-like experiment.  In addition, we have analyzed three effective scenarios of multi-field inflation:  a softly turning, a suddenly turning and a mild waterfall trajectory.  We find that softly turning trajectories do not generate any detectable distortions by PIXIE (they are nevertheless detectable by a PRISM-class experiment, but only for a very limited range of the parameters).  In the case of a sudden turn in the field space, a peak in the scalar power spectrum as well as damped oscillations are predicted on scales exiting the Hubble radius close to the turning point.  We find that an important and potentially detectable level of spectral distortions can be produced if the turning point is tuned to lie within the range of scales relevant for distortions, i.e. about  $8 \\Mpc^{-1}\\lesssim k \\lesssim 10^4 \\Mpc^{-1}$.   A peak in the scalar power spectrum is also a generic prediction of hybrid models ending with a mild waterfall.  The peak can be located in the range of scales that is of the present interest when inflation continues for about 40 e-folds of expansion during the waterfall.  This may lead to an important level of CMB distortions and we have determined the parameter range that will be probed with the PIXIE experiment.   It is important to note that the ratio between $\\mu$-type and intermediate $i$-type distortions could be a good discriminator between the models.  Finally, we have considered another possible scenario potentially leading to a strong distortion signal: a massive curvaton whose blue spectrum is subdominant on CMB angular scales and overtakes the inflaton spectrum on smaller distortion scales.   However, we find that the generated level of distortions does not reach the sensitivity PIXIE when we impose the bounds from ultra compact minihaloes.   But a not yet constrained part of the parameter space could be tested by an ultimate experiment like PRISM.  Our analysis demonstrates that if $\\mu$-type or intermediate $i$-type distortions are discovered by future experiments, they might be produced by only a few inflation models, in particular within multi-field scenarios, with the price being a somewhat tuned and contrived choice of parameters.  Therefore, one might then favour other potential sources of distortions like energy injection from decaying or annihilating particles or from primordial black hole evaporation."
    },
    "1402/1402.7214_arXiv.txt": {
        "abstract": "Water in outflows from protostars originates either as a result of gas-phase synthesis from atomic oxygen at $T$ $\\ga$ 200 K, or from sputtered ice mantles containing water ice. We aim to quantify the contribution of the two mechanisms that lead to water in outflows, by comparing observations of gas-phase water to methanol (a grain surface product) towards three low-mass protostars in NGC1333. In doing so, we also quantify the amount of methanol destroyed in outflows. To do this, we make use of \\textit{JCMT} and \\textit{Herschel}-HIFI data of H$_2$O, CH$_3$OH and CO emission lines and compare them to \\textsc{RADEX} non-LTE excitation simulations. We find up to one order of magnitude decrease in the column density ratio of CH$_3$OH over H$_2$O as the velocity increases in the line wings up to $\\sim$15 \\kms. An independent decrease in $X({\\rm CH_3OH})$ with respect to CO of up to one order of magnitude is also found in these objects. We conclude that gas-phase formation of H$_2$O must be active at high velocities (above 10 \\kms\\, relative to the source velocity) to re-form the water destroyed during sputtering. In addition, the transition from sputtered water at low velocities to formed water at high velocities must be gradual. We place an upper limit of two orders of magnitude on the destruction of methanol by sputtering effects. ",
        "introduction": "\\label{sec:intro} Water plays a unique role in probing the physical and chemical conditions of star-forming regions. In the cold regions of protostars, water is frozen out as ice covering the dust grains, with only trace amounts found in the gas phase \\citep[$\\la$10$^{-7}$;][Schmalzl et al. in prep.]{Caselli2012,Herpin2012,Mottram2013}. The molecular outflows launched by the accreting protostar form an interesting laboratory for studying water in protostellar systems: the high water abundances observed in outflows \\citep[$>$ 10$^{-6}$;][]{Nisini2010,Vasta2012, Santangelo2012, Bjerkeli2012, Dionatos2013,VanLoo2013} arise from a combination of sputtering of ice mantles and direct gas-phase synthesis, however the relative water abundance contributions are unknown. Constraining the relative contributions of these two routes is important for using water as a diagnostic of the physical conditions in outflows, as well as constraining the chemistry in the various parts of the outflow.  Water formation on dust grains is a complicated process \\citep[e.g.][]{Ioppolo2008,Lamberts2013}, but generally involves the repeated hydrogenation of oxygen-bearing species such as O, OH and O$_2$ \\citep{Tielens1982}. The only efficient gas-phase formation route for H$_2$O is a high-temperature ($T \\ga 200$ K) neutral-neutral reaction chain \\citep[e.g.][]{Draine1983, Bergin1999, Tielens2005, Glassgold2009}: \\begin{equation} {\\rm O \\xrightarrow{H_2} OH \\xrightarrow{H_2} H_2O + H}. \\end{equation}  Another common grain-surface material, methanol, can be formed on dust grains through through repeated hydrogenation of CO \\citep[e.g.][]{Tielens1982,Tielens2005} \\begin{equation} {\\rm CO_{surface} \\xrightarrow{H} HCO_{surface} \\xrightarrow{H} CH_3O_{surface} \\xrightarrow{H} CH_3OH_{surface}}. \\end{equation} Unlike water, gas-phase reactions that produce CH$_3$OH are very slow compared to grain surface chemistry for any physical conditions expected within protostellar systems \\citep{Geppert2006}.  Both water \\citep[e.g.][]{Codella2010,vanDishoeck2011} and methanol \\citep[e.g.][]{Bachiller1995,Kristensen2010} have been detected in the gas phase of the outflows and shocks of young stellar objects (YSOs). In astrophysical environments the grain surface reaction chain leads to grain surface abundances of order $10^{-6}-10^{-5}$ for CH$_3$OH with respect to gas-phase H$_2$ \\citep[e.g.][]{VanderTak2000, Cuppen2009}. This results in relative abundances of CH$_3$OH ice with respect to H$_2$O ice ranging from a few to $\\sim$ 30 \\% \\citep[e.g.][] {Dartois1999,Pontoppidan2003,Pontoppidan2004,Gibb2004,Boogert2008,Oberg2011}. Models of cold dense clouds by \\citet{LeeH.-H.1996}, which only include gas-phase production, result in gas-phase abundances of CH$_3$OH no higher than between $10^{-10}$ and $10^{-9}$ on timescales of up to $10^6$ years \\citep{Garrod2006,Geppert2006}.  In the context of YSOs and their outflows, we assume and expect all of the grain-surface formation of H$_2$O and CH$_3$OH to have already taken place during the preceding dark-cloud and prestellar core phases. Material is liberated from the grain surfaces either through thermal desorption if the temperature of the dust grains rises above $\\sim$100 K for water \\citep[][]{Fraser2001, Brown2007} and $\\sim$85\\,K for methanol \\citep[][]{Brown2007} as they journey close to the protostar \\citep[][]{Visser2009}, or through sputtering in shocks. While photodesorption due to cosmic-ray induced UV radiation is expected to be important in the cold conditions found in protostellar envelopes \\citep[e.g.][Schmalzl et al. in prep.]{Caselli2012,Mottram2013}, it is unlikely to be dominant in the conditions found in outflows.  Therefore, any gas-phase CH$_3$OH observed in outflows will originate from sputtered ice mantle material, while gas-phase H$_2$O may originate either in warm shocked gas in outflows where the temperature exceeds $\\sim$ 200 K \\citep{Bergin1999, Charnley1999} or from the grain mantles. \\emph{The first objective of this study is to test whether there is observational evidence in outflow material for water formation in excess of what we would expect to see from only ice mantle desorption, i.e. is the gas phase formation route activated?}  Our method for studying this is to examine the relative gas-phase column density variations of CH$_3$OH and H$_2$O derived in components of their emission spectra associated with the outflow and its shocked cavities, as seen towards the source positions of three prototypical YSOs. CH$_3$OH is a particularly good molecule for comparing against H$_2$O as it is a pure grain mantle species. Thus, any variations of $N({\\rm CH_3OH})/N({\\rm H_2O})$ are a result of independent variations in either $N({\\rm CH_3OH})$ or $N({\\rm H_2O})$.  Before we can use $N({\\rm CH_3OH})/N({\\rm H_2O})$ in studying the variations of $N({\\rm H_2O})$, we must first consider whether any CH$_3$OH is destroyed, either by dissociative desorption during sputtering or in the shock through reactions with H. This latter scenario would be somewhat similar to that involving another grain-surface product, NH$_3$, at positions offset from the YSO \\citep{Codella2010,Viti2011}. \\emph{The second objective of this paper is to provide observational limits on the extent of CH$_3$OH destruction in the outflows of low-mass YSOs.} This is accomplished by comparing the column densities to that of a CO line which traces the same gas as H$_2$O, such as CO J=10$-$9 \\citep{Yldz2013,Tafalla2013,Santangelo2013,Nisini2013}. The $N({\\rm CH_3OH})/N({\\rm CO})$ ratio can then be translated into the CH$_3$OH abundance $X({\\rm CH_3OH}) = N({\\rm CH_3OH})/N({\\rm H_2})$ by assuming a constant CO abundance in the outflow, $10^{-4}$.  This paper is structured as follows. In Section \\ref{sec:observations} we describe our observations and then describe the immediate results in Section \\ref{sec:resuls}. In Section \\ref{sec:analysis} we describe the decomposition of the acquired spectra into components, and explain which of these we assume to have the same physical origin. We also discuss the likely physical conditions in the outflow, and use them to constrain the radiative transfer simulations which allow us to convert our observed line intensity ratios into column density ratios. In Section \\ref{sec:discussion} we discuss the acquired column density ratios and place them in the context of the arguments and processes outlined above. Finally, in Section \\ref{sec:conclusions} we summarise our results and draw our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We combined observations of H$_2$O, CH$_3$OH and CO line emission with \\textsc{RADEX} simulations to constrain the gas-phase abundance variations of water and methanol towards three low-mass YSOs: IRAS 2A, IRAS 4A and IRAS 4B. We find that:  \\begin{enumerate}  \\item The CH$_3$OH/CO column density ratio decreases by up to one order of magnitude with increasing velocity. \\item The abundance of CH$_3$OH in the shocked gas is more than 90$\\%$ lower than that reported for CH$_3$OH ice in the cold envelope. Given that the CO abundance is unaffected by shock chemistry across the velocity range investigated, this implies that CH$_3$OH is destroyed either during the ice sputtering process from dust grains or through gas-phase reactions with H in the outflow. \\item The H$_2$O/CO column density ratios increase by $\\sim$ a factor of two, with the increase beginning at velocities above $\\sim$10 \\kms. This suggests that the gas-phase formation of water is active and significant at velocities higher than 10 \\kms. No discontinuity is observed in the water column density as a function of velocity, implying that the transition is smooth and continuous between the observed water originating from ice sputtering and gas-phase routes. \\item The column density ratio of CH$_3$OH/H$_2$O also decreases with increasing velocity, closely matching the same trend in CH$_3$OH/CO column density ratio. Based on our prior conclusions, this implies that both H$_2$O and CH$_3$OH are sputtered from ices in these shocks, but that only CH$_3$OH is being destroyed, and no gas-phase H$_2$O formation is occurring, at least at velocities below 10 \\kms. \\item In the blue wing of IRAS 4A, the CH$_3$OH/H$_2$O column density ratio decreases more steeply at velocities in excess of $\\sim$10 \\kms\\, than the CH$_3$OH/CO column density ratio. This traces the high-temperature gas-phase formation of water in this higher velocity regime.  \\end{enumerate}  Consequently, gas-phase CH$_3$OH and H$_2$O abundances in shocked regions depend on the complex interplay between ice sputtering mechanisms, gas-phase destruction processes and high-temperature formation reactions. These conclusions hint towards future observational and experimental requirements to further constrain the physics and chemistry of H$_{2}$O and CH$_{3}$OH in outflows. In particular, trying to observe and quantify the efficiency of CH$_3$OH and H$_2$O ice sputtering by neutrals in a laboratory environment would contribute towards verifying and/or refuting two of our main conclusions; whether or not the two molecules are sputtered with equal efficiency, and if CH$_3$OH can be destroyed in the sputtering processes. Observations resolving the outflow spatially close to the protostar, and mapping the variation of physical conditions, are also important to further constrain these results."
    },
    "1402/1402.1354_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small In the core accretion scenario for the formation of planetary rocky cores, the first step toward planet formation is the growth of dust grains into larger and larger aggregates and eventually planetesimals. Although dust grains are thought to grow from the submicron sizes typical of interstellar dust to micron size particles in the dense regions of molecular clouds and cores, the growth from micron size particles to pebbles and kilometre size bodies must occur in the high densities reached in the mid-plane of protoplanetary disks. This critical step in the formation of planetary systems is the last stage of solids evolution that can be observed directly in young extrasolar systems before the appearance of large planetary-size bodies.  \\medskip  Tracing the properties of dust in the disk mid-plane, where the bulk of the material for planet formation resides, requires sensitive observations at long wavelengths (sub-mm through cm waves). At these wavelengths, the observed emission can be related to the dust opacity, which in turns depend on to the grain size distribution. In recent years the upgrade of the existing (sub-)mm arrays, the start of ALMA Early Science operations and the upgrade of the VLA have  significantly improved the observational constraints on models of dust evolution in protoplanetary disks. Laboratory experiments and numerical simulations led to a substantial improvement in the understanding of the physical processes of grain-grain collisions, which are the foundation for the models of dust evolution in disks.  \\medskip  In this chapter we review the constraints on the physics of grain-grain collisions as they have emerged from laboratory experiments and numerical computations. We then review the current theoretical understanding of the global processes governing the evolution of solids in protoplanetary disks, including dust settling, growth, and radial transport. The predicted observational signatures of these processes are summarized.  \\medskip  We discuss the recent developments in the study of grain growth in molecular cloud cores and in collapsing envelopes of protostars as these likely provide the initial conditions for the dust in protoplanetary disks. We then discuss the current observational evidence for the growth of grains in young protoplanetary disks from millimeter surveys, as well as the very recent evidence of radial variations of the dust properties in disks. We also include a brief discussion of the constraints on the small end of the grain size distribution and on dust settling as derived from optical, near-, and mid-IR observations. The observations are discussed in the context of global dust evolution models, in particular we focus on the emerging evidence for a very efficient early growth of grains in disks and the radial distribution of maximum grain sizes as the result of growth barriers in disks.  We will also highlight the limits of the current models of dust evolution in disks including the need to slow the radial drift of grains to overcome the migration/fragmentation barrier. \\\\~\\\\~\\\\~}% ",
        "introduction": "} \\label{Sec:Intro}   In this chapter we will discuss the evolution of dust in protoplanetary disks, focusing on the processes of grain growth and the observational consequences of this process. In the standard scenario for planet formation, this is the phase in which the solids grow from micron-size particles, which are present in the molecular cloud cores out of which stars and protoplanetary disks are formed, to centimeter size and beyond on the path to become planetesimals. This is the last stage of solid growth that is directly observable before the formation of large, planetary-size bodies that can be individually observed. As this phase of growth is directly observable, it has the potential of setting strong constraints on the initial stages of the planet formation process. In this review, we focus on the growth of particles on the disk mid-plane, as this is where most of the solid mass of the disk is concentrated and where planets are expected to form. Grain growth in the disk inner regions and atmosphere can be effectively investigated in the infrared and have been extensively reviewed in \\citet{2007prpl.conf..767N}, although at the time of that review the observational evidence and theoretical understanding of the global dust evolution processes in the disk were limited.  \\medskip  This phase of growth up to  centimeter size particles in protoplanetary disks can also be connected to the study of the most pristine solids in our own Solar System. We can directly follow grain growth in the  Solar nebula through the study of primitive rocky meteorites known as chondrites. Chondrites are composed mainly of millimeter-sized, spherical chondrules with an admixture of smaller, irregularly-shaped calcium aluminum inclusions (CAI) embedded in a fine-grained matrix \\citep[e.g.][]{2005ApJ...623..571S}. Some chondrules may be the splash from colliding planetesimals but most have properties that are consistent with being flash-heated agglomerates of the micron-sized silicate dust grains in the matrix. CAIs are the first solids to condense from a hot, low pressure gas of solar composition.  Their ages provide the zero point for cosmochemical timescales and the astronomical age of the Sun, 4.567 Gyr. Relative ages can be tracked by the decay of short-lived radionucleides and allow the formation timescales of dust agglomerates and planetesimals to be placed in context with astronomical observations \\citep{2011AREPS..39..351D}. Recent measurements of  absolute ages of the chondrules found that these were formed over a period that span from CAI formation to a few Myr beyond \\citep{2012Sci...338..651C}. As chondrule dating refers to the time when the dust agglomerate was melted, this new result tells us  either that dust grains rapidly grew up to millimeter sizes in the early solar nebula and were melted progressively over time while agglomeration into larger bodies continued over the lifetime of the protoplanetary disk; or that the growth of these millimeter and centimeter size agglomerates from smaller particles continued over a period of several million years. In any case, the process of assembling CAIs and chondrules into larger bodies had to occur throughout the disk lifetime (see also the chapter by \\textit{Johansen et al.}).  \\medskip  In this review we will thus concentrate on the most recent developments of the theoretical models for grain evolution in disks and the constraints on the grain-grain collisions from laboratory experiments as well as on the new wealth of data at (sub-)millimeter and centimeter wavelengths that is becoming available. The goal will be to give an overview of the current astronomical constraints on the process and timescale of grain growth in disks.  \\medskip  The processes and observations that we discuss in this chapter are expected to occur in protoplanetary disks and both the theoretical models and the interpretation of the observations rely on assumptions on the disk structure and its evolution. The structure and evolution of protoplanetary disks during the pre-main sequence phase of stellar evolution has been discussed extensively in recent reviews and books \\citep[e.g.][]{2007prpl.conf..555D,2009apsf.book.....H,2010apf..book.....A,2011ARA&A..49...67W}. Throughout this chapter, our discussion will assume a flared, irradiated disk structure in hydrostatic equilibrium with a constant gas to dust ratio of 100 by mass, unless explicitly stated otherwise. While the detailed disk structure and its evolution under the effects of viscous accretion, chemical evolution, photoevaporation and planet-disk interaction are actively investigated in detail (see e.g. the chapters by \\textit{Alexander et al.}, \\textit{Audard et al.}, \\textit{Dutrey et al.}, and \\textit{Turner et al.}), this assumed disk structure is an adequate representation of the early phases of disk evolution and a good reference for the processes of dust evolution.  Figure~\\ref{fig:sketch} shows a sketch of a protoplanetary disk where we illustrate pictorially the physical processes involved in grain growth on the left side. On the right side we illustrate the regions probed by the various observational techniques and the angular resolutions offered by the forthcoming generation of facilities and instruments in the infrared and submillimetre regions.  \\begin{figure*}[thb] \\centering \\includegraphics[width=18cm]{ltesti_f1.pdf} \\caption{Illustration of the structure, grain evolution processes and observational constraints for protoplanetary disks. On the left side we show the main grain transport and collision mechanism properties. The different lengths of the arrows illustrate the different velocities of the different grains. On the right hand side, we show the areas of the disk that can be probed by the various techniques. The axis shows the logarithmic radial distance from the central star. The horizontal bars show the highest angular resolutions (left edge of the bars) that can be achieved with a set of upcoming facilities and instruments for at the typical distance of the nearest star forming regions. } \\label{fig:sketch} \\end{figure*}  \\medskip  In \\S~\\ref{Sec:Processes}, we introduce the main concepts at the basis of the transport, dynamics and evolution of solid particles in disks. First we introduce the interactions between the solids and gas and the complex dynamics of solids in disks, then we describe the basic processes of the growth of solids towards larger aggregates. In \\S~\\ref{Sec:Lab} we discuss the constraint on the outcomes of grain-grain collisions from laboratory experiments and numerical computations. These provide the basic constraints to the global evolution models of solids in disks, which we describe in \\S~\\ref{Sec:Theo}. In \\S~\\ref{Sec:Opacities} we discuss the effects of growth on the dust opacities and the observational consequences. The possible recent evidence for grain evolution at and before the disk formation epoch is discussed in \\S~\\ref{Sec:pre-disc}. The observational constraints for grain growth at infrared wavelengths are described in \\S~\\ref{Sec:scattir}. We discuss grain properties in the disk mid-plane from (sub-)millimeter observations in \\S~\\ref{Sec:Obs}, beginning with the methodology in \\S~\\ref{Sec:Method}, then the results of low resolution, multi-wavelength continuum surveys in \\S~\\ref{Sec:sub-mm_phot}, ending with the most recent resolved studies in \\S~\\ref{Sec:Resolved}. ",
        "conclusions": "Global models of grain evolution in disks, constrained by the results of laboratory and numerical calculations of grain and aggregate collisions, predict that in the conditions of protoplanetary disks grains should very rapidly grow to centimeter sizes. The growth process is not without difficulties, with numerous ``barriers'' to overcome. At this time, the most critical problem seems to be the drift-fragmentation ``barrier'' caused by the large differential radial speeds of grains of different sizes induced by the aerodynamical drag. While the level of growth predicted by models is consistent with observational data, the time evolution is too fast unless large grains radial drift motion is slowed down in pressure ``traps''. Several ideas of plausible mechanisms to produce these traps exist: from disk instabilities to snowlines. The predictions should now be tested observationally and ALMA will offer a unique opportunity to do this in the coming years thanks to its superb angular resolution and sensitivity.  \\medskip  Observations show that the majority of disks around isolated pre-main sequence stars in the nearby star forming regions contain a significant amount of grains grown to at least millimeter sizes. No clear evolutionary trends have emerged so far, although some indirect evidence of disk aging is starting to emerge also from millimeter observations. This finding strongly supports the notion that grain growth is a common and very fast process in disks, possibly occurring in the earliest phases of disk formation. The other implication is that the large grains are either kept in the disk mid-plane  or they are continuously reformed for several Ma. This is required to explain the presence of these dust aggregates in disks around pre-main sequence stars and also consistent with cosmochemical evidence in our own Solar System.  \\medskip  Multi-wavelength radially resolved observations of grain properties, which were just barely becoming possible at the time of the last Protostars and Planets conference, are now providing a wealth of new constraints to the models. The overall radial segregation of dust aggregates of different sizes predicted by models is in general agreement with observations, however, the fast draining of solids towards the inner disk predicted by models cannot reproduce the observations. A mechanism to slow down the radial drift is definitely required.  \\medskip  While the interpretation of the millimeter observations in terms of grain growth appears to be solid, the values of the dust opacities for the dust aggregates are still very uncertain. New computations or laboratory measurements for a range of grain composition and structure would represent a major step to put on very solid grounds the study of dust evolution in disks. Similarly, a major step in the understanding of the physical processes of grain growth in disks will be the extension of the grain-grain laboratory experiments to icy dust particles. These experiments are now been carried out and the initial results are expected in the coming years.  \\medskip  Armed with better constraints on the grain-grain collision outcomes, the global models of dust evolution in disks will also have to be advanced to the next stage. This will have to include a proper account of the gas and dust coevolution and the extension of the models to two and three dimensions, to account for vertical transport and azimuthal inhomogeneities, which are now routinely observed with ALMA.  \\medskip  On the observational side new and upgraded observing facilities at millimeter and centimeter wavelengths are offering an unprecedented opportunity to expand the detailed studies to faint objects and to resolve the detailed radial and vertical structure of the grain properties. In the coming years we expect that it will be possible to put  strong constraints on the evolutionary timescale for the dust on the disk mid-plane and explore this process as a function of the central host star parameters and as a function of environment. The study of the dust and gas small scale structures in the disk, as well as their global distribution, will most likely allow us to solve the long standing problem of radial drift.  \\textbf{ Acknowledgments.} We thank M.~Bizzarro, P.~Caselli, C.~Chandler, C.~Dullemond, A.~Dutrey, Th.~Henning, A.~Johansen, C.~Ormel, R.~Waters, and especially the referee, S. Okuzumi, for valuable comments and discussions on the content of this chapter. We thank F.~Windmark for providing data for Fig.~\\ref{fig:lab}. The work reported here has been partly supported over the years by grants to INAF-Osservatorio Astrofisico di Arcetri from the MIUR, PRIN-INAF and ASI. Extended support from the Munich-IMPRS as well as ESO Studentship/Internship, DGDF and Scientific Visitor programmes in the period 2009-2013 is gratefully acknowledged. S.~Andrews and T.~Birnstiel acknowledge support from NASA Origins of Solar Systems grant NNX12AJ04G. A.~Isella and J. M.~Carpenter acknowledge support from NSF award AST-1109334. J. P. Williams acknowledges support from NSF award AST-1208911.  \\bigskip"
    },
    "1402/1402.3398_arXiv.txt": {
        "abstract": "Planned instruments utilising multi-object adaptive optics systems on the forthcoming extremely large telescopes require large numbers of high order deformable mirrors.  These devices are a significant cost driver, particularly if specifications regarding the number of faulty actuators are stringent.  Here, we investigate the effect on adaptive optics performance that such faulty actuators have, and draw conclusions about how far faulty actuator specifications (and hence cost) can be relaxed without having a significant effect on adaptive optics performance.  We also provide performance estimates using a map of faulty actuators from an existing deformable mirror.  We investigate the effect of faulty actuators using an end-to-end Monte Carlo adaptive optics simulation code.  We find that for actuators stuck at a fixed height above the deformable mirror surface, between 1--2\\% of actuators can be faulty before significant performance degradation occurs.  For actuators that a coupled to nearest neighbours, up to about 5\\%, can be faulty before \\ao performance begins to be affected. ",
        "introduction": "The forthcoming generation of optical ground-based \\elts will have primary mirror diameters of over 30~m.  These facilities will depend on advanced \\ao \\citep{adaptiveoptics} to meet the majority of their scientific goals, and will provide astronomers with the necessary resolutions and light collecting areas to probe the universe with unprecedented sensitivity.  Instruments using advanced modes of \\ao operation are proposed, including \\moao and \\mcao which will deliver \\ao corrected \\psfs over a wide field-of-view.  These systems require more than one \\dm, and \\moao systems in particular require a large number of \\dms, one for each corrected science channel.  Designs for the \\eelt \\eagle instrument \\citep{2010SPIE.7736E..25Rshort} include 20 separately corrected \\ao channels, thus requiring 20 \\dms.  The cost of \\dms are significant factor for consideration in the design of \\moao systems, and any reduction in cost would be welcomed by the design teams.  The failure rate of individual \\dm actuators can also be a risk to these instruments.  Here we investigate the effect of the presence of faulty \\dm actuators on \\moao performance using a full end-to-end Monte-Carlo simulation code, \\dasp \\citep{basden5}.  We model an instrument similar to \\eagle, and consider actuators stuck at different surface heights, including at the full range of the \\dms, and also consider actuators that are coupled to nearest neighbours.  In \\S2, we describe our simulation models and assumptions.  In \\S3, we present our results and findings, and we conclude in \\S4. ",
        "conclusions": "We have made a study of the effect of faulty \\dm actuators on \\elt-scale \\moao performance using an end-to-end Monte-Carlo simulation tool.  Individual \\dm actuators that are stuck at some level relative to the mean surface position have been considered, as well as actuators which are coupled to nearest neighbours.  For the instrument description considered here, we find that between 1--2\\% of actuators can be faulty (stuck at mid-range) before significant performance degradation (of more than a 1--2\\% drop in ensquared energy) occurs.  When actuators are stuck at the stroke limit of the \\dm, the requirements are more stringent, with only 0.2\\% of actuators allowed to be faulty before ensquared energy drops by more than 1\\%.  We find that care should be taken when specifying maximum \\dm stroke if it is known that faulty (stuck) actuators will be present.  For actuators that a coupled to nearest neighbours, a greater fraction, up to about 5\\%, can be faulty before \\ao performance begins to be affected.  It should be noted that more stringent requirements will be needed for instruments demanding higher \\ao performance, such as \\ltao or \\xao systems.  These findings allow the risk and cost for an instrument such as \\eagle to be significantly reduced, and can be applied to any instrument requiring a large number of high order \\dms that are currently at the cutting edge of technological innovation."
    },
    "1402/1402.3167_arXiv.txt": {
        "abstract": "In this work, we show that when supernova Ia (SN Ia) data sets are used to put constraints on the free parameters of inhomogeneous models, certain extra information regarding the light-curve fitter used in the supernovae Ia luminosity fluxes processing should be taken into account. We found that the size of the void as well as other parameters of these models might be suffering extra degenerations or additional systematic errors due to the fitter. A recent proposal to relieve the tension between the results from Planck satellite and SNe Ia is re-analyzed in the framework of these subjects. ",
        "introduction": "According to a homogeneous and isotropic Friedmann-Robertson-Walker (FRW) standard model, since 1998 combined observations of nearby and distant type Ia supernovae (SNe Ia) led to the discovery of the accelerating universe picture. We now have a \\emph{concordance model} in which the dimming of distant SNe Ia \\cite{perlmutter, riess, mik, hicken, kessler, union2, snls, union21, sdss14, sdss14b}, anisotropies in the cosmic microwave background (CMB) \\cite{wmap9, planck}, and the signature of baryon acoustic oscillation (BAO) \\cite{eisen, percival} cannot be explained by considering only baryonic and dark matter. The most popular solution is to introduce an additional component with negative pressure, the so-called dark energy (e.g., \\cite{ht, sahni, padma, frieman, huterer}).  Some other proposals have been presented since then. Among them, exact inhomogeneous models with no dark energy component were put forward shortly after the release of the first supernova measurements \\cite{dabro, pascual, celerier00, tomita}; and more recently some models, such as the ones based on spherically symmetric Lema\\^{\\i}tre-Tolman-Bondi (LTB) and other exact solutions, began to have an important development in the past few years (see for example, \\cite{celerier7, enqvist, garcia, stephon, biswas, bolejko1, clarkson, celerier1373, wessel}). Until today, these have been considered toy-models because of their simplicity and it is necessary to remark that they are not robust models of our universe yet. According to some authors (e.g. \\cite{gbh, moss}), LTB models face important observational challenges and the most simple current versions of these models would be ruled out. As it has been emphasized in the literature, their use must be considered as a mere first step towards more sophisticated models \\cite{celerier1373}. Examples somewhat more complex are the ones known as Swiss-cheese models \\cite{brou}, meatball models \\cite{kain} and Szekeres Swiss-cheese models \\cite{bolecele}. Note that LTB model breaks the Copernican principle, by placing the observer at the centre of a spherically symmetric universe, while other models, such as the Swiss-cheese one, are only locally inhomogeneous. For a detailed review of exact solutions, see \\cite{bolejko1}.  The importance of the study of aforementioned models is motivated by different observational works. In particular, \\cite{keenan} concluded that local measurements of the near-infrared luminosity density are consistent with models that invoke a large local under-density (around 300 Mpc) to explain either the apparent acceleration observed via type Ia supernovae, or to explain the discrepancy between local measurements of $H_0$ and those inferred from the CMB. The use of inhomogeneous models is also well justified even in the standard paradigm. For example, it has already been demonstrated that although the cosmological observations are analyzed in the homogeneous framework, matter inhomogeneities might be mistaken for evolving dark energy \\cite{bolejko2, wessel2}. Other authors have studied how the presence of a local spherically symmetric inhomogeneity can affect apparent cosmological observables derived from the luminosity distance. Under the assumption that the real space-time is exactly homogeneous and isotropic, they have found that phantom dark energy or quintessence behaviors can be produced for compensated underdense or overdense regions \\cite{staro}. The fact of putting observational constraints on these type of models should not be taken to be ruled out or not, but as a beacon to follow, and to study possible degenerations present that might influence future works with more refined proposals.  The possible tension between the best fits for $\\Omega_m$ and $H_0$ obtained from the \\emph{Planck} satellite observations on one hand, and the Hubble diagram of SNe Ia on the other hand, has been recently faced by the authors in \\cite{uzan}. They showed that the use of an inhomogeneous Swiss-cheese model to interpret the Hubble diagram allows to reconcile it with \\emph{Planck} results.  The flux measurements from an SN Ia at different epochs and distinct passbands are processed with the so-called \\emph{light-curve fitters} to obtain luminosity distance values. The two most used methods are named MLCS2k2 \\cite{mlcs} (hereafter MLCS) and SALT2 \\cite{salt2}. Distance moduli calculated for the same objects by the two fitting methods are not necessarily equal.  MLCS (The Multicolor Light Curve Shape fitter), is the most recent version of the fitter used by the High-z Supernova Team \\cite{riess}, whilst the SALT2 (Spectral Adaptive Light curve Template), is an improved version of the fitter used originally by the Supernova Cosmology Project \\cite{perlmutter}. A detailed description of both fitters and a thorough discussion about systematic errors in SN surveys can be found for example in \\cite{hicken, kessler}.  It is a known fact that the same SN Ia data set from which distance estimates are analyzed with two different light-curve fitters, can lead to different values for various cosmological parameters, or also some cosmological models would be more favored than others (e.g. \\cite{hicken, kessler, sollerman, sanchez, carneiro, smale, zhen}). One of us, has recently shown how these two light-curve fitters employed for the same SN Ia data set produce the same result than two different SN Ia sets, and it should be minded as an additional factor to decide whether phantom type models are favored or not \\cite{grb}.  Whereas the MLCS calibration uses a nearby training set of SNe Ia assuming a close to linear Hubble law, SALT2 uses the whole data set to calibrate empirical light curve parameters. SNe Ia beyond the range in which the Hubble law is linear are used, so a cosmological model must be assumed in the latter. Typically a $\\Lambda$CDM or a $w$CDM ($w=const$) model is assumed. Consequently, the published values of SN Ia distance moduli obtained with SALT2 fitter retain a degree of model dependence. Regarding this issue, in \\cite{frieman1} it was pointed out that systematic errors in the method of SNe Ia distance estimation have come into sharper focus as a limiting factor in SN cosmology. The major systematic concerns for supernova distance measurements are errors in correcting for host-galaxy extinction and uncertainties in the intrinsic colors of supernovae, luminosity evolution, U-band rest frame in the low-redshift sample and selection bias. Also, SALT2 fitter does not provide a cosmology-independent distance estimate for each supernova, since some parameters in the calibration process are determined in a simultaneous fit with cosmological parameters to the Hubble diagram. It is important to remark that a 0.2 apparent magnitude difference leads to a 10\\% error in the luminosity distance value. Some researchers have focused on these issues, and important steps have begun to be taken (e.g. \\cite{sullivan, sullivan2, march, lampe, chot, saltmu, gupta, joha, kessler2, wang, wang2, riga, mosher, sdss14b}). Unfortunately, these subjects have been boarded, mostly, just from the SALT2 point of view. Another recently developed light-curve fitter is SiFTO \\cite{sifto}. Although SiFTO differs from SALT2 in some aspects including improvements with respect to the latter, SiFTO shares more features with SALT2 than with MLCS (see for example \\cite{snls}). In fact, results from SiFTO are finally mostly compared using SALT2 as a guide, and the general conclusion is that the differences associated with these two fitters are not very significant (e.g \\cite{sifto, snls}).  Since in the inhomogeneous framework literature the possible effects of the mentioned issues have been scarcely studied; and stimulated by the approach given by \\cite{uzan}, in this Letter, we first analyze the consistency of two of the main light-curve fitters used for the elaboration of SN Ia data sets in the void models framework. To accomplish this, we present a study about the possible results that can be obtained regarding, for instance, the size of the void; or degenerations in other parameters that are usually used in this type of models, to parameterize their profiles or diagnose when viewed as effective models. We also explore the impact of some systematics induced by the mentioned fitters in such models, to then study how this could affect previous proposals in the literature.  Additionally, we propose that the solution found by the authors in \\cite{uzan} might be even more reinforced if the most recent SNe Ia data were used in the MLCS fitter framework. ",
        "conclusions": "In the FRW framework where the universe is isotropic and homogeneous, this is going through an accelerated stage because of the existence of what we call dark energy. Although the evidence seems to be solid from various observational data sets, the search of other alternatives which also explain these observations have been in development in the last few years. That is the case of the so-called inhomogeneous models. Even though these are toy-models yet and do not accomplish the description of the observations correctly, several authors have remarked the importance of their study, and even in the standard paradigm, some degree of inhomogeneity might have a detectable effect on certain observable quantities.  The luminosity distance measurements of SNe Ia constitute the most used data sets to put observational constraints on cosmological models, since they have been and still are the most solid evidence of the acceleration of the universe detected in the framework of a FRW model. But as it is well known in the standard paradigm, when the same SNe Ia data set is processed with two different light-curve fitters (i.e. SALT2 and MLCS), the values found for cosmological parameters (such as the equation of state $w$ of dark energy) differ.  In this work, we analyzed the aforementioned difference showing that, similarly to what occurs in the standard model, the light-curve fitters lead to incompatibilities when SNe Ia data are used to put constraints on inhomogeneous models. This can be seen, for instance, in the $H_0-\\Omega_{\\rm in}$ plane in Fig. \\ref{Fig1}.  We found that when the luminosity fluxes coming from supernovae are processed with the MLCS fitter, the luminosity distances inferred imply sizes of voids 11\\% smaller than in the SALT2 case. The difference found is of the same order that what other authors obtained \\cite{february} when combining SNe Ia data with data of $H(z)$. The fitters seem to have a degeneration with these observations. We also showed that SALT2 favors larger voids and lower $\\Omega_{\\rm in}$.  Fig. \\ref{Fig1} shows evidence that MLCS favors values of $H_0$ lower than in the SALT2 case (something that also occurs in the FRW model). This lead us to analyze the proposal of the authors of \\cite{uzan} to relieve the tension in the $H_0-\\Omega_{m}$ plane between the values of the \\emph{Planck} Collaboration and the one from SNLS 3 SNe Ia data \\cite{snls}. Although in our work, when using SDSSII data \\cite{kessler}, the mentioned tension with \\emph{Planck} does not appear, maybe the solution found by those authors might be even more reinforced if the most recent SNe Ia from SNLS data were used in the MLCS fitter framework. We would find interesting the release of more public data sets processed with MLCS to be used in the framework of models alternative to $\\Lambda$CDM. Note that other authors have warned about the risk of using tables with luminosity distance values in the framework of SALT2 to put constraints on alternative models (e.g. \\cite{smale, frieman1}). Given that MLCS constitutes a more model-independent fitter than SALT2, data processed with the former fitter should be preferentially chosen to put constraints to alternative models when using SN Ia data. Alternatively, it would be interesting to have luminosity distances tables in the framework of proposals like the one developed in \\cite{saltmu}.  We found that MLCS tends to favor an effective deceleration today ($q_{\\rm eff}>0$) with more emphasis that the SALT2 case. Other authors have already found these trends, but in the framework of dark energy models \\cite{zhen}.  When analyzing cases in which the value of $H_0$ is fixed (for example, to see what would happen if one would want to make the mentioned value compatible with the one obtained from CMB), as it is expected, the allowed ranges for the size of the void are reduced. But, in the same situation (fixing $H_0$), both fitters lead to very different results although the very same data has been used. Also, for MLCS the ranges associated with the size of the void get more restricted than in the case of SALT2. It is interesting to highlight that the size of the void for MLCS, under this situation, might even be smaller than 1 Gpc; indicating that when using MLCS, the size of the voids might not need to be so giant as it is usually sustained.   \\bigskip"
    },
    "1402/1402.2027_arXiv.txt": {
        "abstract": "{The origin of the Halley-type comets (HTCs) is one of the last mysteries of the dynamical evolution of the Solar System. Prior investigation into their origin has focused on two source regions: the Oort cloud and the Scattered Disc. From the former it has been difficult to reproduce the non-isotropic, prograde skew in the inclination distribution of the observed HTCs without invoking a multi-component Oort cloud model and specific fading of the comets. The Scattered Disc origin fares better but suffers from needing an order of magnitude more mass than is currently advocated by theory and observations. Here we revisit the Oort cloud origin and include cometary fading. Our observational sample stems from the JPL catalogue. We only keep comets discovered and observed after 1950  but place no a priori restriction on the maximum perihelion distance of observational completeness. We then numerically evolve half a million comets from the Oort cloud through the realm of the giant planets and keep track of their number of perihelion passages with perihelion distance $q<2.5$~AU, below which the activity is supposed to increase considerably. We can simultaneously fit the HTC inclination and semi-major axis distribution very well with a power law fading function of the form $m^{-k}$, where $m$ is the number of perihelion passages with $q<2.5$~AU and $k$ is the fading index. We match both the inclination and semi-major axis distributions when $k \\sim 1$ and the maximum imposed perihelion distance of the observed sample is $q_{\\rm m} \\sim 1.8$~AU. The value of $k$ is higher than the one obtained for the Long-Period Comets (LPCs), for which typically $k \\sim 0.7$. This increase in $k$ is most likely the result of cometary surface processes. We argue the HTC sample is now most likely complete for $q_{\\rm m}<1.8$~AU. We calculate that the steady-state number of active HTCs with diameter $D>2.3$~km and $q<1.8$~AU is of the order of 100.} ",
        "introduction": "The Solar System is host to a large population of comets, which tend to be concentrated in two components, i.e. the Oort cloud and Scattered Disc \\citep{1997Sci...276.1670D}. When examining the orbital data of new comets entering the inner solar system, \\cite{1950BAN....11...91O} noted that the distribution of reciprocal semi-major axis of these comets showed a sharp increase when $1/a<5 \\times 10^{-4}$~AU$^{-1}$. The observed semi-major axis distribution led Oort to suggest that the Sun is surrounded by a cloud of comets in the region between 20\\,000~AU to 150\\,000~AU, and that it contains approximately $10^{11}$ comets with isotropic inclination and random perihelia. Oort showed that the action of passing stars is able to account for the isotropic inclination distribution, random perihelia, the continuous injection of new comets on trajectories that lead to perihelia close enough to the Sun to become visible, and the fact that most new comets have near-parabolic orbits. This hypothesized cloud of comets surrounding the Sun is now called the 'Oort cloud' (OC). \\\\ The only method with which we can infer the number of comets in the OC is by determining the flux of long-period comets (LPCs) and Oort cloud comets (OCCs). The former are comets with a period longer than 200~yr that have had multiple passages through the solar system while the latter are comets that are having their first perihelion passage in the inner solar system.\\\\ For a new comet to enter the planetary region from the OC, it must have sufficiently low angular momentum to allow it to penetrate to small perihelion distances. For a given comet, its angular momentum per unit mass is given by $v_{\\rm t} r$, where $v_{\\rm t}$ is its transverse velocity and $r$ is the distance from the Sun. In the OC, where $r$ is large, $v_{\\rm t}$ must be very low to allow the comet to eventually come close to the Sun. Oort (1950) recognized this and added that the maximum allowable $v_{\\rm t}$ for a given distance $r$ would be a function of $r$. He gave this the name `loss cone' because comets within the velocity cone could enter the planetary region and be ejected or captured to a shorter period orbit by the perturbations of the giant planets, principally Jupiter and Saturn (which would require a perihelion distance $q<15$ AU). In other words, the comets would be `lost' from the OC. These `lost' comets remain in the loss cone until ejected or destroyed. Thus the presence of Jupiter and Saturn presents an obstacle for new comets to overcome in order to enter the inner solar system. This is often called the Jupiter-Saturn barrier.\\\\ There exist two agents which perturb the comets in the cloud onto orbits that enter the inner solar system: passing stars \\citep{1980A&A....85..191W, 1981AJ.....86.1730H} and the Galactic tide \\citep{1986Icar...65...13H, Levison2001}. The passing stars cause usually small random deviations in the orbital energy and other orbital elements. The Galactic tide on the other hand systematically modifies the angular momentum of the comets at constant orbital energy. \\cite{1986Icar...65...13H} discovered that if the semi-major axis of the comet is large enough then the comet's change in perihelion can exceed 10~AU in a single orbit and 'jump' across the orbits of Jupiter and Saturn and thus not suffer their perturbations before becoming visible; in this particular case the comet is considered a new comet. However, the \\cite{1986Icar...65...13H} approximation of the Galactic tide only used the vertical component, which is an order of magnitude stronger than the radial components. In this approximation the $z$-component of the comet's orbital angular momentum in the Galactic plane is conserved and the comets follow closed trajectories in the $q-\\omega$ plane (with $q$ being the perihelion distance and $\\omega$ being the argument of perihelion). Including the radial tides breaks this conservation and the flux of comets to the inner solar system from the OC is increased \\citep{Levison2006}. The trajectories that lead comets into the inner solar system should be quickly depleted, were it not for the passing stars to refill them \\citep{2008CeMDA.102..111R}. The synergy between these two perturbing agents ensures there is a roughly steady supply of comets entering the inner solar system. The story is slightly more complicated because not all comets need to jump over the Jupiter-Saturn barrier: a small subset is able to sneak through and undergo small enough energy changes that they still formally reside in the cloud \\citep{2010MNRAS.404.1886K, 2013Icar..222...20F}, but the main argument still holds. \\\\ The OCCs form just a subset of comets that are visible from the Oort cloud. The OCCs are generally thought to be on their first passage through the inner solar system, but may have had a few passages beyond $\\sim$5~AU. However, we also have observational evidence of comets that are dynamically evolved. These are the LPCs: they have undergone many passages through the realm of the giant planets and/or the inner solar system and most of them are dynamically decoupled from the Oort cloud. A subset of these are the Halley-type comets, or HTCs, named after its progenitor, comet 1P/Halley.\\\\ The second reservoir of cometary objects is the Scattered Disc \\citep{1997Sci...276.1670D}. This source population resides on dynamically unstable orbits beyond the orbit of Neptune. These bodies dynamically interact with Neptune and reside generally on low-inclination orbits. The Scattered Disc is the most likely source of the so-called Jupiter-family comets (JFCs) \\citep{1997Sci...276.1670D, 2008ApJ...687..714V, 2013Icar..225...40B}.\\\\ Visible comets tend to historically be categorized into several groups: the JFCs, HTCs, LPCs and OCCs. Here we adhere to the definitions of \\cite{1996ASPC..107..173L} with a few modifications. These are: \\begin{itemize} \\item JFCs: $T_{\\rm J} \\in (2,3)$ and $a<7.35$~AU (Period $P<20$~yr), \\item HTCs: $T_{\\rm J} <2$, $q>0.01$~AU and $a < 34.2$~AU (Period $P<200$~yr), \\item LPCs: $T_{\\rm J}<2$, $q>0.01$~AU and  $34.2<a<10\\,000$~AU (Period $P>200$~yr), \\item OCCs: $q>0.01$~AU and $a>10\\,000$~AU. \\end{itemize} Here $a$ is the semi-major axis, $P$ is the period of a comet, and $T_{\\rm J}$ is the Tisserand parameter with respect to Jupiter. The criterion $q>0.01$~AU was applied in order to exclude the Sun-grazing Kreutz comets. Dormant comets were also excluded in our work. We realize that the above populations are not all inclusive: we did not include the Encke type and the Chiron type. There is no definition for comets with $T_{\\rm J}>2$ and $P>20$~yr, though these are generally included in the Centaur population.\\\\ The origin of the HTCs has long been controversial. The best attempts to determine the source of the HTCs were done by \\cite{1994A&A...281..911F}, \\cite{Levison2001}, \\cite{2002MNRAS.333..835N}, and \\cite{Levison2006}.\\\\ \\cite{1994A&A...281..911F} investigated the dynamical behaviour of LPCs coming from the Oort cloud and the role of perturbations by the planets. They used a combination of a Monte Carlo random walk method and \\\"{O}pik's equations \\citep{1951PRIA...54..165O} and discovered that an initially isotropic distribution may result in a flattened distribution of evolved comets because of  the inclination dependence on planetary energy kicks and the tendency of near-perpendicular comets to become retrograde. They find that the typical comet undergoes 400 revolutions before becoming extinct, and combined with their capture probability they suggest there are of the order of 300 active HTCs.\\\\ \\cite{Levison2001} studied an Oort cloud origin for the HTCs but concluded that the non-isotropy of the HTCs could only be reproduced with an Oort cloud that had a non-isotropic structure: it had to be flattened in the inner parts and isotropic in its outer parts. \\cite{Levison2001} included cometary fading in their simulations but they found that this alone was not enough to match the HTC inclination distribution. \\cite{2002MNRAS.333..835N} reached a different conclusion based on results of their Monte Carlo simulations: the outer Oort cloud could be the source of the HTCs and their number and inclination distribution could be fairly well reproduced when cometary fading and disruption was taken into account. However, when adding the inner Oort cloud they produced far too many HTCs, which \\cite{Levison2001} suggested was needed in their model. Unfortunately \\cite{2002MNRAS.333..835N} did not try to match the semi-major axis distribution, something that Levison et al. (2001) did try to do. In addition, the initial conditions of both of these works were artificial because they started the comets on planet-crossing orbits.\\\\ These issues led \\cite{Levison2006} to suggest that the HTCs originate from the outer edge of the Scattered Disc. This latter reservoir has a flattened inclination distribution \\citep{1997Icar..127...13L} and the effects of the Galactic tide at the outer edge is able to supply the retrograde HTCs. Their simulations match the inclination and semi-major axis distribution of the HTCs reasonably well. The problem with this scenario is that it requires the Scattered Disc to be more than an order of magnitude more massive than most estimates thus far \\citep{Gomes2008, 2013Icar..225...40B}. This leads us back to the Oort cloud being the most likely possible source.\\\\ Here we aim to demonstrate that the Oort cloud is the most likely source of the HTCs and that a single dynamical model with cometary fading is sufficient to reproduce their current inclination and semi-major axis distributions. This paper is organized as follows.\\\\ In the next section we present the observational dataset that we use for comparison. Section~3 details our numerical methods. In Section~4 we present the results of our simulations and comparison with the observed dataset followed by a short discussion in Section~5. We present a summary and conclusions in the last section. ",
        "conclusions": "In this work we have used numerical simulations to determine whether the distribution of the HTCs can be reproduced using the Oort cloud as their source distribution and when considering cometary fading. We generated a large number Oort cloud comets that were evolved numerically under the gravitational influence of the Sun, giant planets, Galactic tide and passing stars using well-established methods. Guided by the work of \\cite{Wiegert1999} we post-processed the data with a power law fading mechanism that was applied to all comets that passed perihelion closer than 2.5~AU. We compared the numerical output to the JPL catalogue, keeping only those comets discovered and observed after 1950 to limit observational biases.\\\\ We find that the Oort cloud with cometary fading is the most likely source of the HTCs. Both of their semi-major axis and inclination distributions can be well reproduced when the fading index $k\\sim1$ and when one considers only those comets for which $q \\lesssim 1.8$~AU. The HTC fading index is somewhat higher than the value of 0.6 we find for LPCs. This increase in fading index could be the result of physical processes on the cometary surface. With these results we estimate there are of the order of 100 active HTCs with diameters $D>2.3$~km and $q<1.8$~AU at any time."
    },
    "1402/1402.2161_arXiv.txt": {
        "abstract": "\\noindent Probing the evolution of the universe at high redshifts with standard candles is a powerful way to discriminate dark energy models, where an open question nowadays is whether this component is constant or evolves with time. One possible source of ambiguity in this kind of analyses comes from cosmic opacity, which can mimick a dark enery behaviour. However, most tests of cosmic opacity have been restricted to the redshift range  $z<2$.  In this work, by using luminosity distances of gamma-ray bursts (GRBs), { given the validity of the Amati relation}, and the latest $H(z)$ data we determine constraints on the cosmic opacity at high redshifts ($z>2$) for a flat $\\Lambda$CDM model. A possible degenerescence  of the results with the adopted cosmological model is also investigated by considering a flat XCDM model. The limits on cosmic opacity in the redshift range $0<z<2$ are updated with type Ia supernovae (SNe Ia) from the Union2.1 sample, where we added the most distant ($z=1.713$) spectroscopically confirmed SNe Ia. From the analyses performed, we find that both SNe Ia and GRBs samples  are compatible with a transparent universe at $1\\sigma$ level and the results are independent of the dark energy equation of state parameter $w$. ",
        "introduction": "\\label{sec:introduction}  Type Ia supernovae (SNe Ia) observations provide the most direct evidence for the current cosmic acceleration. In the context of Einstein's general theory of relativity, this result implies either the existence of some sort of dark energy, constant or that varies slowly with time and space (see \\cite{caldwell,li} for recent reviews), or that the matter content of the universe is subject to dissipative processes \\cite{Lima1999,chimento}. On the other hand, since SNe  Ia observations are affected by at least four different sources of opacity, namely, the Milky Way, the hosting galaxy,  intervening galaxies, and the intergalactic  medium, alternative mechanisms contributing to the acceleration evidence or even mimicking the dark energy behaviour have been proposed. Examples are possible evolutionary effects  in SNe Ia events \\cite{drell,combes}, a local Hubble bubble \\cite{zehavi,conley}, modified gravity \\cite{Ishak,ks2007,Bertschinger},  unclustered sources of light attenuation \\cite{Aguirre,Rowan,Goobar} and the existence of Axion-Like-Particles (ALPs), arising in a wide range of well-motivated high-energy physics scenarios, and that could lead to dimming of SNe Ia brightness (\\cite{Av2}, for a review of the phenomenology of the weakly-interacting-sub-eV-particles (WISPs) see \\cite{Jaeckel}).  In the last years, the  cosmic distance duality (CDD) relation has been used to verify the existence of exotic physics as well as the presence of opacity and systematic errors in SNe Ia observations. The CDD relation  connects the luminosity distance, $D_{\\scriptstyle L}$, to the angular diameter distance, $D_{\\scriptstyle A}$, by \\cite{eth33,ellis07} \\begin{equation} \\frac{D_{\\scriptstyle L}}{D_{\\scriptstyle A}}{(1+z)}^{-2}=1. \\label{rec} \\end{equation} This relation is valid for all cosmological models based on Riemannian geometry, requiring only that source and observer are connected by null geodesics in a Riemannian spacetime, that the number of photons is conserved \\cite{ellis07} and narrow-beam effects are negligible \\cite{busti2012,clarkson}.  Several methods have been proposed in the literature to test modifications of the CDD relation. For instance, \\cite{bk04} used SNe Ia data as measurements of $D_{\\scriptstyle L}$  and estimates of $D_{\\scriptstyle A}$ from FRIIb radio galaxies \\cite{daly} and ultra-compact radio sources \\cite{G94} to test possible deviations of the CDD relation. They found a 2$\\sigma$ violation caused by an excess brightening of SNe Ia at $z > 0.5$, perhaps due to lensing  magnification bias. \\cite{h3,h2,h3a} proposed a test of the CDD relation by using luminosity distances of SNe Ia  and angular diameter distances of galaxy clusters. The verification of the CDD relation validity depended on the assumptions used to describe the galaxy clusters and the light-curve fitters  used in SNe Ia analysis (see also \\cite{Nair,xi,nam}).  Still with SNe Ia and angular diameter distances of galaxy clusters data,  \\cite{li2013} explored the impact of cosmic opacity on the cosmic acceleration observation. Their results suggested that an  accelerated cosmic expansion is still needed to account for the dimming of SNe Ia and the standard cosmological scenario remains to be supported by current observations. By using four different parameterizations to possible CDD deformations and SNe Ia data \\cite{Amanullah} the authors of ref. \\cite{Lima} found that only if the cosmic opacity is fully negligible,  the description of an accelerating universe powered by dark energy or some alternative gravity theory must be invoked. { In Ref. \\cite{HAC} measurements of the expansion rate $H(z)$ and Union 2.1 SNe Ia compilation were used to impose cosmological model-independent constraints on cosmic opacity \\cite{foot}. It was found that a completely transparent universe  is in agreement with the largest SNe Ia sample.}  More recently, an interesting test of the CDD relation via cosmic microwave background spectrum was proposed by \\cite{ellis2013}. As a result, it was found that the CDD relation cannot  be violated by more than 0.01\\% between the decoupling era and today.  Furthemore, a cosmological model-independent test involving only measurements of the gas mass fraction of galaxy clusters from Sunyaev-Zeldovich and X-ray surface brightness observations was discussed by \\cite{h5b}, where no significant  violation of the CDD relation was found, supporting a transparent universe in the X-ray and radio wavelengths.   In face of the discussed above, most tests of cosmic opacity  have been limited in the redshift range $0 <  z < 1.7$. On the other hand, since current SNe Ia data even when combined with all other currently available data sets cannot yet determine whether the energy density of dark energy is constant as required by the $\\Lambda$CDM model or time-varying, some authors have used luminosity distances of gamma-ray bursts (GRBs) to explore the universe at higher redshifts ($2<z<8$) to understand how dark energy behaves. GRBs provide a means to test cosmological models once phenomenological relations are used to transform them into stardardizable candles (see \\cite{amati,ghirlanda}). For example, they were used to test dark energy models \\cite{schaefer2007,wang,ratra2010,wei2} and also possible bias from underdense lines of sight \\cite{busti_prd_2012}. In this way, in order to have confidence in analyses involving GRBs it is important to verify the existence of cosmic opacity also in high redshifts.  In this paper, we  probe cosmic opacity at higher redshifts through luminosity distances from GRBs data, { by assuming the validity of the Amati relation \\cite{amati}}, and with the latest 26 opacity-free  measurements of the Hubble expansion in the  range ($0.1<z<2.30$) using the approach proposed by Avgoustidis {\\it et al.} \\cite{Av,Av2} (see next section to details). We find that GRB data are in a very good agreement with a transparent universe, with bounds consistent with the ones derived with SNe Ia data. A possible degenerescence of the opacity constraints with the underlying cosmological model is also investigated, where it is shown that the opacity parameter $\\epsilon$ is nearly insensitive of a different dark energy equation of state parameter $w$ for a flat XCDM model.   The paper is organized as follows. In Sec. II we describe the methodology. In Sec. III  the observational data we use in the statistical analyses are presented. The corresponding constraints on the cosmological opacity  are investigated in Sec. IV. The article is finished with a summary of the the main results in the conclusion section. ",
        "conclusions": "In the last few years, several cosmological observables, such as baryon accoustic oscillations, luminosity distances of SNe Ia, angular diameter distances of galaxy clusters, gas mass fraction and $H(z)$ measurements have been used to probe cosmic opacity and/or searching for some evidence of new physics. However, due to limitations in the redshift distribution of the data, most of these analyses were limited in the redshift range $0<z<2$, where the region between the SNe Ia and the cosmic microwave background data remained unexplored.  In this work, we have probed cosmic opacity in the the redshift range $1.5<z<8$ through luminosity distances from GRBs, { by assuming the validity of the Amati relation \\cite{amati}}, and the latest 26 measurements of the Hubble expansion from passively evolving galaxies and baryon acoustic oscillations in the range ($0.1<z<2.30$). In order to test a dependence of the method with the adopted cosmology, we  compared the constraints from flat $\\Lambda$CDM model with the flat XCDM model. We found that GRB data are in full agreement with a perfect transparent universe and the results are independent of $w$. We parameterized the cosmic opacity such as $\\tau(z)=\\epsilon z$. By marginalizing over $\\Omega_M$, we obtained $\\epsilon=0.06\\pm 0.18$ $(1\\sigma)$ (flat $\\Lambda$CDM) and, by marginalizing over $\\Omega_M$ and $w$, we obtained $\\epsilon=0.057\\pm 0.21$ at $1\\sigma$ level for a flat XCDM model. We also used the Union2.1 SNe Ia sample, where we added the most distant ($z=1.713$) spectroscopically confirmed SNe Ia to impose limits on opacity in redshift range $0<z<2$. From the joint analyses involving SNe Ia and $H(z)$ data we got $\\epsilon=0.020 \\pm 0.055$ $(1\\sigma)$ (flat $\\Lambda$CDM) and $\\epsilon=0.015\\pm 0.060$ $(1\\sigma)$ (flat XCDM). We believe that the results presented here reinforce the interest in the observational search for GRBs. When larger samples with smaller statistical and systematic uncertainties become available one may improve the limits to cosmic opacity at high redshifts ($z>2$) as well as test new parameterizations for $\\tau(z)$, which can help us to pinpoint the nature of dark energy.    \\label{sec:conclusions}"
    },
    "1402/1402.7111.txt": {
        "abstract": "We develop a high-precision model for laser ranging interferometric (LRI) observables of the GRACE Follow-On (GRACE-FO) mission. For this, we study the propagation of an electromagnetic wave in the gravitational field in the vicinity of an extended body, in the post-Newtonian approximation of the general theory of relativity. We present a general relativistic model for the phase of a plane wave that accounts for contributions of all the multipoles of the gravitating body, its angular momentum, as well as the contribution of tidal fields produced by external sources. We develop a new approach to model a coherent signal transmission in the gravitational field of the solar system that relies on a relativistic treatment of the phase. We use this approach to describe high-precision interferometric measurements on GRACE-FO and formulate the key LRI observables, namely the phase and phase rate of a coherent laser link between the two spacecraft. We develop a relativistic model for the LRI-enabled range between the two GRACE-FO spacecraft, accurate to less than 1~nm, and a high-precision model for the corresponding range rate, accurate to better than 0.1~nm/s. We also formulate high-precision relativistic models for the double one-way range (DOWR) and DOWR-enabled  range rate observables originally used on GRACE and now studied for interferometric measurements on GRACE-FO.  Our formulation justifies the basic assumptions behind the design of the GRACE-FO mission and highlights the importance of achieving nearly circular and nearly identical orbits for the GRACE-FO spacecraft. ",
        "introduction": "\\label{sec:intro}  The Gravity Recovery and Climate Experiment (GRACE) was a very successful 2002 US-German space mission. A pair of satellites spent nine years mapping the gravitational field of the Earth using a highly accurate microwave ranging system between the two spacecraft, which were flying in nearly identical orbits. GRACE demonstrated the feasibility of monitoring temporal variations in the Earth's gravitational field and thus detect, for instance, both seasonal variations and long term changes in the hydrosphere.  GRACE's planned successor, tentatively called the GRACE Follow-On (GRACE-FO) mission, is scheduled for launch in 2017. In addition to a microwave ranging system similar to that of GRACE and operating in the Ka-band, GRACE-FO will be equipped with a laser ranging interferometer (LRI) instrument \\cite{Pierce-etal:2008}, a system using heterodyne optical interferometry. The LRI is expected to provide range with an accuracy of 1~nm and corresponding estimates for range rate.  The LRI experiment on board GRACE-FO will rely on two-way measurements, designating one spacecraft as master and the other as transponder (we shall call them GRACE-FO-{\\rm A} and GRACE-FO-{\\rm B}, correspondingly) and utilizing a phase-locked active laser transponder with a frequency offset \\cite{Sheard_etal:2012}. Phase locking in the transponder is expected to eliminate transponder laser frequency noise, leaving the transmitter laser frequency noise and phase noise due to pointing errors, experimental features that are captured in the instrument design \\cite{Pierce-etal:2008,Sheard_etal:2012}, as the two major noise sources.  Although the original GRACE model formulation was based on Newtonian arguments (see details in \\cite{Kim:2000}), from the early stages of the GRACE-FO mission development it became clear that at the level of accuracy expected from GRACE-FO, general relativistic effects may become significant. The two-way nature of the experiment reduces the contribution of some of these effects to the GRACE-FO observables, yet their influence on the modeling of these quantities must be carefully analyzed. The mathematical model of the ultimate LRI observables---time series data obtained at the master spacecraft, recording continuous changes in the phase difference between the local laser oscillator and the laser beam returned by the transponder spacecraft---must take into account general relativistic contributions to the length of the signal path and differences between coordinate time and proper time.  In this paper, we focus on the formulation of a relativistic model for computing and processing the observables of the GRACE-FO mission. We rely on a previously developed theory of relativistic proper reference frames within a system of $N$ extended bodies and the motion of light and test particles in the vicinity of an extended body \\cite{Turyshev-Toth:2013}. The organization of the paper is as follows: In Sec.~\\ref{sec:em-phase} we discuss light propagation in the vicinity of the extended Earth and derive a general relativistic solution for the phase of an electromagnetic wave. In Sec.~\\ref{sec:LRI-obs} we discuss the process of forming the inter-satellite LRI observables of GRACE-FO and derive a model for the phase and the relevant LRI-enabled range observable. We also develop a relativistic model for the frequency observable and a related model for interferometric range rate. We develop equations to model observable fluctuations in the phase rate and range acceleration. In Sec.~\\ref{sec:dowr} we discuss the possibility of extracting dual-one-way (DOWR) style observables from the interferometric data, similar to that used by the GRACE-FO microwave ranging system. We develop the relevant equations and models for range, range-rate and range acceleration observations. We conclude with a set of recommendations and an outlook in Sec.~\\ref{sec:sonc}.  In an attempt to streamline the presentation of results and to keep the main body of the paper focused, we present some relevant details in the form of appendices. In Appendix~\\ref{sec:GR-background} we discuss the post-Newtonian approximation of general relativity, introduce the metric tensor in this formulation, and discuss its properties. In Appendix \\ref{sec:ranges} we present details of the derivation of instantaneous distances between the spacecraft. In Appendix \\ref{sec:need-for-J2} we emphasize the need to include properly the quadrupole moment of the Earth into relativistic coordinate transformations, equations of light propagation and equations of the motion of geocentric satellites. Finally, in Appendix \\ref{sec:useful-rels} we introduce some useful relations that help in the evaluation of the magnitudes of various expressions that involve combinations of orbital parameters of the GRACE-FO spacecraft.  The notational conventions employed in this paper are those used in %by Landau and Lifshitz \\cite{Landau-Lifshitz:1988}. Letters from the second half of the Latin alphabet, $m, n,...=0...3$ denote spacetime indices. Greek letters $\\alpha, \\beta,...=1...3$ denote spatial indices. The metric $\\gamma_{mn}$ is that of Minkowski spacetime with $\\gamma_{mn}={\\rm diag}(+1,-1,-1,-1)$ in the Cartesian representation. We employ the Einstein summation convention with indices being lowered or raised using $\\gamma_{mn}$. We use powers of $G$ and negative powers of $c$ as bookkeeping devices for order terms. Other notations are explained as they occur. ",
        "conclusions": "\\label{sec:sonc}  We studied, in the post-Newtonian approximation of the general theory of relativity, the propagation of a plane electromagnetic wave traveling through the gravitational field in the vicinity of an extended body. An arbitrarily shaped and rotating distribution of matter is represented by an arbitrary energy-momentum tensor. We derived a compact closed form general relativistic solution describing the phase of a plane wave that accounts for contributions of all the mass and current multipoles of the body itself (Eqs.~(\\ref{eq:eik-sol-all}) and (\\ref{eq:eik-sol-all-S})), as well as the contribution due to the tidal gravity field (given by Eq.~(\\ref{eq:eq_eik-phi-lamb-T1})) produced by the external bodies forming a system of $N$ astronomical bodies. As such, the solution that we obtained significantly extends previous similar derivations found in the literature.  We evaluated a solution for relativistic phase (\\ref{eq:phase_t}), which, in addition to the usual Shapiro term, accounts for the quadrupole term of the Earth's mass distribution, the Earth's spin, and the tidal gravity introduced by the potentials of external bodies in the GCRS.  At the level of accuracy anticipated from GRACE-FO, the recommended solution for relativistic phase is given by (\\ref{eq:eik-phi-lamb-EJ}). This formulation allows one to achieve a self-consistent analytical result that can be evaluated for the GRACE-FO LRI experiment. All the necessary information regarding the geometry of the experiment and the background gravitational field is captured in the total signal path ${\\cal R}$, a quantity we introduced in Sec.~\\ref{sec:GRACE-FO-phase}.  Based on the established solution, we presented a new formulation for the relativistic phase transformation that describes a coherent signal transmission between the two spacecraft of the GRACE-FO constellation. We developed LRI observables for GRACE-FO including both the phase and phase rate of the signal received at the master spacecraft, including relativistic treatment of the transponder offset frequency. Eq.~(\\ref{eq:dn_ABA2=+*=6}), together with  the relativistic range model (\\ref{eq:range-ABA0}), provide a high-precision formulation for the LRI phase observable on board GRACE-FO. Similarly, Eq.~(\\ref{eq:range-rate-0}) provides a high-precision model for the phase rate observable on board GRACE-FO. These results will allow GRACE-FO to reach the desired resolution of 1 nm in range and 1~nm/s in range rate.  Our formulation justifies the basic assumptions behind the design of the GRACE-FO mission. In particular, our analysis demonstrates the importance of achieving nearly circular and nearly identical orbits (with eccentricity of $e\\sim 0.001$) for the twin GRACE-FO spacecraft. If these requirements  are satisfied, the observables can be represented by a set of very simple models, allowing the project to streamline data conditioning and ultimate science data analysis. Conversely, should the orbital parameters be less ideal than assumed (as presented in Table~\\ref{tb:params}; see also Appendix~\\ref{sec:useful-rels}), the range and range rate models would have to be updated to include terms that we were able to omit, as their magnitudes were sufficiently small for the planned, ideal orbits. Of particular concern are several omitted terms in the range rate model (\\ref{eq:range-rate}), namely the $1/c^2$ Sagnac term given by (\\ref{eq:term-c-3.1}), the Shapiro term discussed in (\\ref{eq:shap1}), and the quadrupole term, evaluated in (\\ref{eq:j2-eval}). If needed, the omitted terms can be easily identified and reinstated by retracing the computational steps presented in this paper.  The analysis of relativistic effects presented here dealt mainly with the magnitudes of the largest terms of these contributions. These are either constant or vary at the orbital frequency (i.e., 1/rev). Our analysis can be used to validate the choices made in mission design, including decisions concerning terms that are assumed to be removable in the form of an empirically introduced bias. Moreover, our results may be easily extended to analyze the contributions of other periodic terms that come at twice the orbital frequency, i.e., 2/rev, and higher. The analytical models that we present for the range and range rate are, therefore, also applicable in situations that take the mission outside the empirical domain, for instance, when optimal orbital parameters are not achieved and the aforementioned empirical bias can no longer be reliably used.  In addition to developing a model for relativistic contributions to the two-way GRACE-FO LRI observable, we also considered the possibility that the laser-ranging instrument will be used in a DOWR configuration. In this operating mode, the transmitters on the two spacecraft are no longer operating in coherent mode, and the quality of the observable relies on the frequency stability of the two independent laser oscillators on board the twin spacecraft.  As a part of our analysis, we identified an intriguing possibility that may lead to an improvement in the accuracy of the Eddington parameter $\\gamma$. As we discussed in Sec.~\\ref{sec:phasemag}, if the GRACE-FO laser interferometer instrument achieves a 1~nm range resolution, this experiment could improve the accuracy of the estimate of the Eddington's relativistic parameter $\\gamma$ to $5.3\\times 10^{-6}$, which exceeds by a factor of 5 or more the accuracy of the current best estimate for $\\gamma$, provided by the Cassini mission to Saturn \\cite{Bertotti-Iess-Tortora:1999}. This possibility is truly remarkable; it is indicative  of the high level of accuracy achieved by modern geodesy when an engineering team needs to account for a number of general relativistic effects in order to reach the stated science objectives of an Earth-orbiting mission.  The analysis presented in this paper was conducted using an idealized set of conditions without considering noise. It can nonetheless be used to study the propagation of various forms of noise through the GRACE-FO architecture and investigate the impact of various noise sources on future science investigations with GRACE-FO. The formulation for the GRACE-FO LRI observables presented here allows for the direct introduction of noise terms (such as the frequency stability of the on-board laser oscillator or the accuracy of orbit determination), as part of an investigation of the noise contribution to the ultimate accuracy of the experiment.  One of the principal objectives of GRACE-FO is to monitor seasonal changes in the Earth's gravitational field using long-duration data spans. In this analysis, the difference between the orbital frequency and the frequency of thermal noise due to solar heating can be significant, as it can introduce an unmodeled annual term that can mask or enhance estimates of annual variability. Therefore, it is important to model the corresponding effects correctly and disentangle relativistic, thermal, and instrumental effects.  Concluding, we emphasize that the eikonal-based approach and the corresponding solution for the phase of a plane electromagnetic wave developed in Sec.~\\ref{sec:em-phase} provides an efficient way to account for general relativistic effects on the propagation of light in the solar system, as needed, for instance, for the high-precision astrometric campaign recently initiated by the ESA's Gaia mission \\cite{Klioner:2003} (see details at {\\tt http://sci.esa.int/gaia/}). Also, the new approach we presented in Sec.~\\ref{sec:LRI-obs} makes it possible to develop a highly accurate description of signal (re)transmission in the post-Newtonian approximation of a gravitational theory. This approach is based on the generally covariant notion of the phase of a plane wave and allows one to formulate observables in the proper reference frames of the transmitting and receiving spacecraft. The results are applicable to other past and planned missions with similar observables, notably to GRACE and GRAIL missions \\cite{Turyshev:2012nw,Konopliv-etal:2013} and new space gravimetry missions like Satellite-to-Satellite Interferometry (SSI) mission \\cite{Cesare-etal:2006} that is currently studied. Furthermore, our approach can be readily applied in other contexts, including experiments conducted away from the Earth, that involve high-precision measurements between multiple spacecraft relying on precision phase measurements done with microwave signals \\cite{Asmar2005} or laser ranging interferometry (or both): for instance, the formulation of the TDIR observables for the LISA mission \\cite{Tinto:2004yw,Shaddock:2004ua}. Such possibilities will be investigated and the results will be reported elsewhere."
    },
    "1402/1402.2357_arXiv.txt": {
        "abstract": "{ We present recent results from an adaptive optics imaging survey of 40 Luminous IR Galaxies (LIRGs) searching for obscured core collapse supernovae and studying the galaxies themselves.  Here, in particular, we discuss the Super Star Clusters (SSC) populations in the LIRGs.  We have constructed the first statistically significant samples of Luminosity Functions (LF) of SSCs in the near-IR, and find evidence that the LF slopes in LIRGs are shallower than in more quiescent spiral galaxies.  Distance and blending effects were investigated in detail paving the way for SSC studies further out than done previously. We have also correlated the luminosities of the brightest clusters with the star formation rates (SFR) of the hosts. The relation is similar, though somewhat steeper than that found in the optical and at lower SFR levels, suggesting systematic extinction and/or age effects. We find that the characteristics of the relation suggest an underlying physical driver rather than solely a size-of-sample effect.  In particular, a truncated luminosity/mass function would naturally explain the small scatter we find. Finally, we are modelling the ages and masses of our near-IR detected clusters in conjunction with HST optical data and present early results of using SSC properties to trace the histories of the target LIRG systems. } ",
        "introduction": "\\label{sec:intro}  We are conducting a survey of approximately 40 Luminous Infrared Galaxies (LIRGs; galaxies with $10^{11} >  L_{IR}  > 10^{12}$)  using adaptive optics (AO) NIR imaging  with VLT/NACO, Gemini/ALTAIR/NIRI, and recently also using the superb \"wide field\" capability of multi-conjugate AO on Gemini/GEMS.  The instruments deliver images with a spatial resolution perfectly complementing existing HST optical data for many of the targets.  The sample galaxies are at distances ranging from 40 to 180~Mpc, are at various stages of merging, interaction, or isolation and are  also observed with optical spectroscopy at SALT/RSS in Sutherland, South Africa (e.g. \\cite{vaisanen08}).  Local LIRGs might physically represent more closely high-redshift star-formation (SF), where a diversity of modes of strong SF are present (e.g. \\cite{rodighiero11}).  Our survey is dubbed SUperNovae and starBursts in the InfraReD, {\\em SUNBIRD}, for the twofold science aim:  We search for very obscured core-collapse SNe close to the nuclei of star-forming galaxies to determine the total SFR in the local universe (see e.g. \\cite{mattila12,kankare12}), and study the LIRGs themselves to have a uniform sample of local LIRGs to investigate the physical details of star-formation and its triggering and history as a function of interaction stage and type, environment, and metallicity (\\cite{vaisanen08,vaisanen12}).   In these proceedings we report on the galaxies themselves, and in particular on their Super Star Cluster (SSC) populations which likely trace extreme forms of SF in the systems. ",
        "conclusions": ""
    },
    "1402/1402.0493_arXiv.txt": {
        "abstract": "We present an integral field spectroscopic study of the Luminous Infrared Galaxy (LIRG) NGC 7130, a known starburst-AGN composite galaxy. We employ standard emission line ratio diagnostics and maps of velocity dispersion and velocity field to investigate how the dominant ionising sources change as a function of radius. From the signatures of both star formation and AGN activity we show that NGC 7130 is a remarkably clean case of starburst-AGN mixing. We find a smooth transition from AGN dominated emission in the centre to pure star forming activity further out, from which we can estimate the radius of the extended narrow line region to be 1.8 $\\pm$ 0.8 kpc. We calculate that the fraction of \\OIII\\ luminosity due to star formation and AGN activity is 30 $\\pm$ 2\\% and 70 $\\pm$ 3\\% respectively, and that the fraction of \\Ha\\ luminosity due to star formation and AGN activity is 65 $\\pm$ 3\\% and 35 $\\pm$ 2\\% respectively. We conclude with a discussion of the importance and potential of starburst-AGN mixing for future studies of the starburst-AGN connection. ",
        "introduction": "\\label{Sec:intro}  One of the longest standing mysteries in astrophysics is the nature of the relationship between star formation and AGN activity in composite galaxies. For many years, it has been well established that there is a strong relationship between the accretion activity of supermassive black holes and the evolution of their host galaxies \\citep{Cattaneo99, Haehnelt00, Granato04, LaMassa12, Ka03, CidFernandes01}. The mass of a supermassive black hole has been found to scale with properties of the host galaxy such as velocity dispersion (M-$\\sigma$ relation), the stellar mass in the bulge ($M_{BH}$-$M_{\\star}$ relation) and the luminosity of the bulge ($M_{BH}$-L relation) \\citep[e.g.][]{Magorrian98, Marconi03, Bennert11, McConnell13, Ferrarese00, Gebhardt00, Tremaine02, Gultekin09}. Although these relationships indicate that by z = 0, the building of the stellar mass and black hole mass in galaxies is in some way correlated, it is unclear whether such a correlation is the result of a connection between black hole and star formation activity on timescales relevant to the formation of stars (facilitated by processes such as mergers, starburst-driven winds and AGN-driven outflows, e.g. \\citealt{Yuan10, Rafferty11}), or whether these scaling relations are simply pointing to something more fundamental which is yet to be uncovered \\citep[e.g.][(review)]{Mullaney12,Alexander12}.  Naively, one might suggest that the scaling relations simply fall out of the common requirement of cold gas to fuel both star-formation and AGN activity. This explanation is problematic, as gas must lose $\\approx$ 99.9\\% of its angular momentum in order to move from a $\\approx$ 10 kpc stable orbit in a star-forming disk to the zone of influence of an AGN on a $\\approx$ 10 pc scale \\citep{Jogee06}. However, it is well established both theoretically and observationally that black hole activity must impact star formation. Black hole feedback can heat material in the inter-stellar medium of a galaxy \\citep[e.g.][]{Sijacki06}, or drive powerful outflows resulting in significant gas redistribution \\citep[e.g.][]{Alatalo11, Rich10, Rich11, DeBuhr11}. These processes deplete the galaxy's cold gas reservoir, resulting in the regulation or quenching of star formation \\citep[e.g.][]{DiMatteo05, Croton06, Hopkins06, Sijacki07, Hopkins08, Booth09, McCarthy10}. Black hole feedback directly impacts the results of galaxy evolution simulations, which consistently indicate that AGN feedback is pivotal for preventing over-production of stellar mass in galaxies \\citep[e.g.][]{Vogelsberger13, Croton06}.  However, the finding that approximately 30-50\\% of Seyfert 2 nuclei are associated with young stellar populations \\citep[e.g.][]{Sarzi07, GonzalezDelgado05, CidFernandes04, GonzalezDelgado01, Storchi-Bergmann01} seems to suggest that AGN activity may enhance star-formation, and/or that circumnuclear star-formation may enhance AGN activity. \\citet{Silverman09} found that the vast majority of AGN hosts at z $\\leq$ 1 have star formation rates higher than that of non-AGN hosts of similar mass, and that the incidence of AGN activity increases with decreasing stellar age.  All this evidence suggests that star formation itself is an important source of viscosity in the circumnuclear gas, assisting some of the gas to fall into the AGN, while driving the rest into states of higher angular momentum. In order to gain greater insight into the starburst-AGN connection, it is necessary to study the relationship between black hole and star formation activity on relatively instantaneous timescales. However, it is notoriously difficult to accurately calculate quantities such as star formation rates in starburst-AGN composite galaxies. In such galaxies, the photometric and spectroscopic indicators of star formation will be contaminated by emission from the AGN, so that it is only possible to derive upper limits on star formation rates, with an extremely large margin for error.  Several studies have attempted to combat this issue by calculating the relative contribution of AGN and star formation in starburst-AGN composite galaxies. \\citet{Heckman04} compute the AGN fractions for galaxies from the Sloan Digital Sky Survey (SDSS; \\citealt{York00}) Data Release 4 (DR4) by comparing the positions of galaxies on the \\NIIHa\\ vs \\OIIIHb\\ diagnostic diagram to synthetic composite galaxy templates with different AGN fractions. \\citet{Davies07} estimate the contribution of the AGN to the line emission of a sample of Seyfert galaxies by calculating the equivalent width of a CO absorption feature and comparing this to the known value for globular clusters. \\citet{Imanishi11} constrain the contribution of the AGN in the infrared by analysing the strength of polycyclic aromatic hydrocarbon (PAH) features in the 3-20$\\mu$m range. All of these methods assume that AGN fractions can be reliably calibrated across the AGN sequence. However, variations in ionisation parameter, metallicity, hardness of the radiation field and electron density across the active galaxy population introduce significant uncertainty in the derived AGN fractions when information from multiple galaxies is combined (e.g. \\citealt{Ke13}).  The advent of integral field spectrographs has provided astronomers with the tools required to conduct accurate and detailed studies of galaxy power sources. Nuclear spectra, such as those obtained by SDSS, can provide information regarding only the very centre of the galaxy (requiring correction for significant aperture effects, see \\citealt{Ke05}). However, given several hundreds or even thousands of spectra from across the galaxy, it is possible to study the variation of the ionisation state and velocity field of gas throughout the galaxy. To date, some of the most elegant results from IFU data have related to the study of shock-starburst composite galaxies. For example, \\citeauthor{Rich10} (\\citeyear{Rich10,Rich11}) used IFU data to form the basis of a convincing argument for the identification of two Luminous Infrared Galaxies (LIRGs) - IC 1623 and NGC 3256 - as shock-starburst composites, and in another instance to show that the extended LINER-like emission of NGC 839 can be best described by shocks originating from a galactic superwind. \\citet{Fogarty12} also used line ratio maps and velocity field information from the Sydney AAO Multi-Object IFS (SAMI) to identify a galactic wind in ESO 185-G031.  IFU data are also of particular interest to the study of the starburst-AGN connection, as they allow for the separation of star-forming and AGN dominated regions. Already, several IFU studies, both in the optical and in the infrared, have identified and investigated starburst-AGN composite galaxies, such as NGC 1068 \\citep{Gerssen06}, NGC 5135 \\citep{Bedregal09}, IRAS 19254-7245 \\citep{Reunanen07} and HE 2211-3903 \\citep{Scharwachter11}. IFU data also provide the information required to study how the contribution of the starburst and the AGN vary with distance from the central AGN, probing the radius of the AGN Extended Narrow Line Region (ENLR).  In this paper, we capitalise on the the strengths of IFU data to carry out a detailed study of the ionisation sources and kinematics in NGC 7130. NGC 7130 is a LIRG, observed using the Wide Field Spectrograph (WiFeS) on the ANU 2.3m telescope as part of the Great Observatory All-Sky LIRG Survey (GOALS; \\citealt{Armus09}). LIRGs are particularly appropriate targets for study of the starburst-AGN connection as they commonly contain both starbursts and AGNs. The prevalence of dual energy sources in LIRGs is likely to be a direct result of their history as mergers of molecular gas-rich spirals, which are thought to be the trigger for powerful bursts of star formation (e.g \\citealt{Sanders96, Barnes92}). During merger events, gaseous dissipation funnels material towards the centre of mass of coalescing systems, leaving them with extremely large reservoirs of gas. This gas rotates in the gravitational field of the galaxy, and can be converted into a rotating stellar disk by means of a starburst (e.g. \\citealt{Barnes92}, \\citeauthor{Mihos94} \\citeyear{Mihos94,Mihos96}) and/or feed AGN accretion processes.  Evidence for the starburst-AGN composite nature of NGC 7130 has been found in the infrared, optical, UV and X-ray (\\citealt{Spinoglio02, Levenson02, Contini02, GonzalezDelgado98, CidFernandes01, GonzalezDelgado01}). In this paper, we use this galaxy to show that IFU data can be utilised to create spatial maps which allow a clean separation of regions of star formation, AGN and composite ionisation from both sources, and permit the identification of outflows and/or feedback that may be present in the galaxy. We demonstrate that there is a strong relationship between the ionisation state of the gas and distance from the centre of the galaxy - a clear indicator of starburst-AGN mixing. We explore how well IFU data can resolve the size of the AGN extended narrow line region and the overall contribution of the starburst and the AGN to optical line emission in individual spaxels.  We summarize the existing observational framework for the composite starburst-AGN nature of NGC 7130 in Section \\ref{Sec:Sec2}. We discuss our observations and data reduction in Section \\ref{Sec:Sec3} and present emission line maps in Section \\ref{Sec:Sec4}. In Section \\ref{Sec:Sec5} we present emission line diagnostic diagrams, discuss indicators of starburst-AGN mixing and quantify the size of the AGN ENLR, calculate starburst and AGN fractions and investigate how these fractions vary as a function of radius. We analyse velocity field information in Section \\ref{Sec:Sec6}, finding further support for the composite nature of this galaxy. We discuss the implications of our findings in Section \\ref{Sec:Sec7} and present our conclusions in Section \\ref{Sec:Sec8}.  \\begin{figure} \\centerline{\\includegraphics[scale=0.65]{NGC7130_pointing_3.eps}} \\caption{WiFeS Pointing (solid green rectangle) overlaid on a WFPC2 image of NGC 7130 taken with the \\emph{Hubble Space Telescope} (F606W, 6030$\\AA$; \\citealt{Malkan98}).}\\label{Fig1} \\end{figure}  \\begin{figure*} \\centerline{\\includegraphics[scale=0.7]{NGC7130_maps.eps}} \\caption{Top to bottom, left to right: Maps of (a) \\NII/H$\\alpha$, (b) \\SII/H$\\alpha$, (c) \\OI/H$\\alpha$, and (d) \\OIII/H$\\beta$. The maps have dimensions $25''\\times38''$, with the same orientation as the HST image in Figure \\ref{Fig1}. Cyan crosshairs illustrate the location of the AGN (as identified by the \\OIHa\\ peak) on all line ratio maps. The nuclear region of the galaxy clearly has enhanced line ratios, indicating the presence of an AGN. A ring of star-formation is also visible in the \\NIIHa\\ map. }\\label{Fig3} \\end{figure*}  Throughout this paper we adopt cosmological parameters $H_{0} = 70.5{\\rm kms}^{-1}{\\rm Mpc}^{-1}$, ${\\Omega}_{\\Lambda} = 0.73$, and $\\Omega_{M}=0.27$ based on the 5-year WMAP results by \\citet{Hinshaw09} and consistent with flat $\\Lambda$-CDM cosmology. ",
        "conclusions": "\\label{Sec:Sec8}  Using IFU emission line data, we have shown that NGC 7130 is a remarkably clean prototype for starburst-AGN mixing in composite galaxies. By combining line ratio maps and line diagnostic diagrams we have uncovered a clear relationship between the location of points on the starburst-AGN mixing sequence and their distance from the centre of the galaxy. We find that:  \\begin{itemize} \\item{The spatial variation of the strong line ratios in NGC 7130 is consistent with a combination of star-formation and AGN activity.} \\item{The composite spaxels, identified using the \\NIIHa\\ vs \\OIIIHb\\ diagnostic diagram, form a clear ring around the nuclear region of the galaxy.} \\item{Spaxels which lie further up the mixing sequence on the \\NIIHa\\ vs \\OIIIHb\\ diagnostic diagram also lie closer to the centre of the galaxy. When color-coding spaxels according to their distance up the mixing sequence we observe clean rings of gas ionised by different combinations of star formation and AGN activity; indicating that the ionisation state of the gas decreases smoothly with distance from the central AGN.} \\item{We use the clean demarcation between \\HII\\ region and ENLR gas to calculate the size of the AGN ENLR. We find the radius of the ENLR to be 1.8 $\\pm$ 0.8 kpc, consistent with expected values from literature.} \\item{We present an empirical method for calculating the relative contribution of the star-formation and the AGN activity on a spaxel-by-spaxel basis. We use this to determine that the fraction of \\OIII\\ luminosity from star-formation and AGN activity is 30 $\\pm$ 2\\% and 70 $\\pm$ 3\\% respectively, and that the fraction of \\Ha\\ luminosity from star-formation and AGN activity is 65 $\\pm$ 3\\% and 35 $\\pm$ 2\\% respectively. This highlights the degree of error made when using \\Ha\\ as a proxy for star-formation rate in composite galaxies.} \\item{Our investigation of the velocity field in NGC 7130 supports the notion of starburst-AGN composite activity, with three separate velocity components identified - one tracing pure star-formation and the underling velocity field of the galaxy, one tracing starburst-AGN mixing over the extended narrow line region of the galaxy, and one tracing the strongly blueshifted and high dispersion nuclear region of the galaxy.} \\item{The emission line and velocity field maps for NGC 7130 also suggest the presence of an outflowing wind in the galaxy; however, our data do not allow us to determine whether this wind is driven by the circumnuclear starburst or the AGN.} \\end{itemize}  Starburst-AGN mixing has the potential to open the doors to understanding the relationship between star formation and AGN activity, both in individual galaxies and across the entire AGN sequence. In Paper II we will expand this study to a larger sample of galaxies and present starburst-AGN mixing models which provide further support for the observational scenario we have presented. We will also discuss applications of starburst-AGN mixing in more detail, with particular focus on the ability to construct individual spectra for the star-formation and AGN components on a spaxel-by-spaxel basis and the possibility of using observations of starburst-AGN mixing to constrain the conditions of the interstellar medium and the hardness of the AGN ionizing radiation field in starburst-AGN composite galaxies."
    },
    "1402/1402.2519_arXiv.txt": {
        "abstract": "{ Recently, \\citet{moriya2013} developed an analytic bolometric light curve model for supernovae interacting with dense circumstellar media. Because of the dense circumstellar medium, the shocked region is assumed to be radiative and make a thin dense shell. The model is based on the conservation of momentum in the shocked dense shell. However, the analytic model was mentioned to neglect the `snow-plow' phase of the shocked dense shell by \\citet{ofek2014}. The `snow-plow' or momentum-conserving phase refers to the period in which the momentum injection from the supernova ejecta is almost terminated and the radiative shocked dense shell keeps moving due only to the momentum previously provided by the supernova ejecta. In this Note, I clarify that the analytic model of \\citet{moriya2013} does take the `snow-plow' phase into account and the criticism of \\citet{ofek2014} is incorrect. In addition, \\citet{ofek2014} related the sudden luminosity break observed in the light curve of Type IIn SN 2010jl to the transition to the `snow-plow' phase. However, I argue that the sudden transition to the `snow-plow' phase is not consistent with the luminosity break observed in SN 2010jl. The luminosity break is likely to be related to other phenomena like the dense shell exiting the dense part of the circumstellar medium. } ",
        "introduction": "Mass loss of massive stars has critical roles in many aspects of astrophysical phenomena \\citep[e.g.,][]{smith2014}. One of the outstanding problems regarding the stellar mass loss is the existence of supernova (SN) progenitors having extremely high mass-loss rates immediately before their explosion. Especially, Type IIn SNe are known to show the signatures of the dense circumstellar media (CSM) which are presumed to be related to the extreme mass-loss activities of the progenitors immediately before their explosion \\citep[e.g.,][]{fransson2002,taddia2013,kiewe2012}. Recently, \\citet[][M13 hereafter]{moriya2013} developed a bolometric light curve (LC) model for SNe interacting with the dense CSM. The model is used to interpret many Type IIn SN LCs to estimate the general properties of the mass loss of Type IIn SN progenitors \\citep{moriya2014}. Similar LC models are suggested previously for the CSM from the steady mass loss \\citep[e.g.,][]{svirski2012,wood-vasey2004} but M13 considered more general cases. \\citet{ofek2014} have also developed a similar LC model to the M13 model to interpret the LC of SN 2010jl, which is one of the most well-observed Type IIn SNe so far (see \\citealt{fransson2013,ofek2014} and the references therein). One interesting feature in the LC of SN 2010jl is the luminosity break observed at around 320 days after the first detection of the SN \\citep{fransson2013,ofek2014}. \\citet{ofek2014} related the break to the transition of the forward shock propagating in the dense CSM to the `snow-plow' or momentum-conserving phase \\citep{svirski2012}. If the shocked region by the SN explosion is radiative as is the case in Type IIn SNe, the shocked region makes a thin dense shell \\citep[e.g.,][]{chevalier1994}. The `snow-plow' phase corresponds to the phase when the momentum supply from the SN ejecta to the shell is almost terminated. Roughly speaking, the shell enters the `snow-plow' phase when the mass of the shocked dense CSM is comparable to the mass of the SN ejecta. \\citet{ofek2014} mentioned that the `snow-plow' phase is neglected in the model of M13.  In this Note, I will first clarify that the criticism of \\citet{ofek2014} on the M13 LC model regarding the `snow-plow' or momentum-conserving phase is incorrect. I will show that the `snow-plow' phase is properly taken into account in the M13 model. I confirm this by following the evolution of the shocked dense shell numerically. Finally, I discuss the luminosity break observed in SN 2010jl and argue that the sudden luminosity break observed in SN 2010jl is inconsistent with the transition to the `snow-plow' phase. ",
        "conclusions": "I have shown that the analytic LC model of M13 which is developed to model SNe interacting with dense CSM takes the 'snow-plow' phase, or momentum-conserving phase, of the shocked dense shell into account. The model was criticized for neglecting it by \\citet{ofek2014} but the criticism is incorrect. The M13 LC model can reproduce the LC behavior expected in the `snow-plow' phase well ($L\\propto t^{-1.5}$ in the steady wind, see \\citealt{svirski2012}). The analytic LC model is also compared to the numerical model and the analytic model is shown to reproduce the numerical result well.  The luminosity break observed in SN 2010jl was related to the transition to the `snow-plow' phase by \\citet{ofek2014}. However, the sudden luminosity break observed in SN 2010jl is not expected from the transition to the `snow-plow' phase. The transition is expected to occur gradually as the shocked shell accumulates the CSM mass. The sudden transition may occur if there is a high-density shell in the CSM but the collision of the shocked shell to the high-density shell is expected to be accompanied by the luminosity increase, not the break. Thus, the luminosity break is not likely to be related to the transition to the `snow-plow' phase."
    },
    "1402/1402.1340_arXiv.txt": {
        "abstract": "We consider Kaluza-Klein (KK) models where internal spaces are compact Einstein spaces. These spaces are stabilized by background matter (e.g., monopole form-fields). We perturb this background by a compact matter source (e.g., the system of gravitating masses) with the zero pressure in the external/our space and an arbitrary pressure in the internal space. We show that the Einstein equations are compatible only if the matter source is smeared over the internal space and perturbed metric components do not depend on coordinates of extra dimensions. The latter means the absence of KK modes corresponding to the metric fluctuations. Maybe, the absence of KK particles in LHC experiments is explained by such mechanism. ",
        "introduction": "\\setcounter{equation}{0}  One of the main tasks of the  Large Hadron Collider (LHC)  is to search for physical phenomena beyond the standard model. If such events occur, they should suggest the directions of the further development of physics. Extra spatial dimensions are among these phenomena. The idea of multidimensionality of our Universe demanded by the theories of unification of the fundamental interactions is one of the most breathtaking ideas of theoretical physics. It takes its origin from the pioneering papers by Th. Kaluza and O. Klein \\cite{KK}, and now the most self-consistent modern theories of unification such as superstrings, supergravity and M-theory are constructed in spacetimes with extra dimensions (see, e.g., \\cite{Polchinski}). One of the clearest evidence for the existence of extra dimensions is the detection of Kaluza-Klein (KK) modes/particles which correspond to appropriate eigenfunctions of the internal manifold, i.e. excitations of fields in the given internal space. Such excitations were investigated in many articles (see, e.g., the classical papers \\cite{SS,Nieu,Kim}). However, up to now, such particles were not detected in the LHC experiments (see, e.g., \\cite{ATLAS,ATandCMS}). In our letter we present a possible explanation of this problem in the case of KK models where stabilized internal spaces are compact Einstein spaces. We perturb this multidimensional background by gravitating masses and show that multidimensional Einstein equations are consistent only if the gravitating bodies are uniformly smeared over the internal space and, consequently, the metric fluctuation components depend only on the coordinates of the external/our space. Obviously, this means the absence of KK modes corresponding to the metric fluctuations. We also demonstrate that the coupling constant between gravitating masses and radions depends on the equation of state parameter $\\Omega$ for the gravitating masses in the internal space. In the case of black strings/branes ($\\Omega=-1/2$) the coupling disappears and the gravitating masses do not perturb the internal space. This explains the excellent agreement of black strings/branes with the gravitational tests.  \\ ",
        "conclusions": "In this letter, we have considered KK models where internal spaces are compact Einstein spaces, e.g., orbifolds. These spaces are stabilized by background matter, e.g., monopole form-fields. We perturbed this background by a compact matter source with the zero pressure in the external/our space and an arbitrary pressure in the extra dimensions. For example, such matter source can be modeled by a gravitating mass. Then, we investigated the metric perturbations in the weak-field limit and showed that the Einstein equations are compatible only if both the metric perturbations and the energy density of the compact matter source do not depend on the coordinates of the internal space. For gravitating masses, this means their smearing over the internal space. For the metric perturbations, this means the absence of the KK modes. The KK modes/particles are the object of the active search in LHC experiments. However, up to now such modes are not detected \\cite{ATLAS,ATandCMS}. As we show in our letter, the reason for this may be smearing of particles and fields over the internal spaces, although it looks unnatural from the point of view of statistical physics because any nonzero temperature should result in fluctuations, i.e. in KK states.  It should be mentioned that the metric and field perturbations are often considered without taking into account the reason of such fluctuations (see, e.g., \\cite{KLZ,BD}). Our analysis clearly shows that the inclusion of the matter sources, being responsible for the perturbations, imposes strong restrictions on the model. Moreover, these restrictions are not caused directly by the made choice of the energy-momentum tensors of the sources. For instance, KK excitations do exist under a similar choice with $\\Omega=0$ in the case of toroidal compactification \\cite{EZ1}, but then there is a contradiction with the relativistic gravitational tests \\cite{EZ3}.  Besides, it is worth noting that our investigation is related also to the popular Universal Extra Dimension models \\cite{theor bounds UED} if the internal spaces are orbifolds. Clearly, the brane world models require a separate investigation. There are also models with combined KK and brane topology \\cite{RS+KK}. Our research shows that in such combined models the smearing of gravitational sources can also occur.   \\"
    },
    "1402/1402.0078.txt": {
        "abstract": "We present a detailed stellar clustering analysis with the application of the two-point correlation function on distinct young stellar ensembles. Our aim is to understand how stellar systems are assembled at the earliest stages of their formation. Our object of interest is the star-forming region NGC\\,346 in the Small Magellanic Cloud. It is a young stellar system well-revealed from its natal environment, comprising complete samples of pre--main-sequence and upper main-sequence stars, very close to their formation. We apply a comprehensive characterization of the autocorrelation function for both centrally condensed stellar clusters and  self-similar stellar distributions through numerical simulations of stellar ensembles. We interpret the observed autocorrelation function of NGC\\,346 on the basis of these simulations. We find that it can be best explained as the combination of two distinct stellar clustering designs, a centrally concentrated, dominant at the central part of the star-forming region, and an extended self-similar distribution of stars across the complete observed field. The cluster component, similar to non-truncated young star clusters, is determined to have a core radius of $\\sim$\\,2.5\\,pc and a density profile index of $\\sim $\\,2.3. The extended fractal component is found with our simulations to have a fractal dimension of $\\sim$\\,2.3, identical to that found for the interstellar medium, in agreement to hierarchy induced by turbulence. This suggests that the stellar clustering at a time very near to birth behaves in a complex manner. It is the combined result of the star formation process regulated by turbulence and the early dynamical evolution induced by the gravitational potential of condensed stellar clusters. ",
        "introduction": "It is generally accepted that stars form in groups of various sizes and characteristics \\citep{ladalada03}, starting with small compact concentrations of protostars embedded in star-forming regions and moving up in {\\sl length-scale} to large extended loose aggregates of young stars stars. It is, however, suggested that these diverse stellar assembles are not independent from each other, but tightly connected through the star formation process \\citep{elmegreen11}. Small, dense proto-clusters coexist in a symbiotic fashion with larger, less dense subgroups of OB-type stars, which in turn reside in even larger, looser stellar associations, each of these types of objects representing a different {\\sl time-scale} of star formation within one molecular cloud \\citep{efremovelmegreen98}. This picture of clustered star formation at various length- and time-scales is not always clear in our observations, as e.g., in the massive star formation environments of star-burst clusters.  Studies of newly formed stellar systems can identify the conditions that may favor {\\sl multiple} over {\\sl single} cluster formation events. Different scenarios for star formation predict different observable properties for the resulting stellar systems. The `quiescent' star formation scenario \\citep[e.g.][]{krumholztan07} predicts large age-spreads among young stars in the same molecular cloud, while according to the `competitive accretion' scenario \\citep[e.g.][]{clark07} star formation is a process of short time-scale, leading to clusters in a variety of forms.  The investigation of the clustering of stars at the time of their formation can provide important information on the nature of the star formation process itself \\citep[e.g.,][]{schmeja08}, and place constraints to the suggested theories. A single compact stellar cluster, embedded in its own H{\\sc ii} region, would suggest a local monolithic episode of star formation in a dense environment. However, such clusters are rarely found in isolation; they are the densest stellar concentrations of larger stellar structures as seen in dwarf and spiral galaxies \\citep{efremov09, karampelas09}. In our own Galaxy multiple (or fractured) clusterings of stars in large star-forming regions, are also more commonly observed \\citep{feigelson11, megeath12}. This clustering behavior favors the existence of multiple processes, such as feedback, controlling star formation, and being active on different scales.  %% If this stellar clustering multiplicity is also hierarchical %% is still to be investigated.  Giant Molecular Clouds are hierarchical structures \\citep{elmegreen96, stutzki98}, indicating that scale-free processes  determine their global morphology. Turbulence is being widely accepted as the dominant among these processes \\citep{maclowklessen04, elmegreenscalo04}. The goal of this paper is to establish whether the newly formed stars follow a similar, hierarchical distribution, which may indicate that turbulence also determines the clustering behavior of stars at the time of their formation. This investigation requires a complete census of stars, covering a high dynamic range in masses, distributed over the typical length-scale of giant molecular clouds.  %The investigation of structures of young stars would also allow to assess the time-scale within which the assembling of stars in various clusterings occurs. %Our motivation is to understand how stars are grouped in various clusterings, and how these clusterings relate to each other by birth.  Our target of interest is the star-forming complex NGC\\,346, the brightest {\\sc H\\,ii} region in the Small Magellanic Cloud \\citep[LHA\\,115-N66;][]{henize56}. {\\sl Hubble} Space Telescope imaging of such regions in the Magellanic Clouds (MCs) provides an unprecedented access to their newly-born stellar populations down to the sub-solar regime over large areas of the sky. Observed by {\\sl Hubble} (GO Program 10248; PI: A. Nota), this region satisfies, thus, the main observational criteria for our analysis: 1) Observed on size-scales relevant for molecular clouds (50\\,-\\,100\\,pc); 2) High angular resolution ($\\sim$\\,0.125\\arcsec); 3) Large number of detected members, covering a significant fraction of the stellar Initial Mass Function, complete to $\\sim$\\,0.5 M{\\solar}. %Moreover, MCs have little line-of-sight extinction, and are less affected by confusion from fore- and background stars than Galactic regions. Our dataset, obtained with the Advanced Camera for Surveys (ACS), is described in \\cite{gouliermis06}.  A rich sample of more than 98,000 stars was detected down to $m_{\\rm 555} \\simeq 27$~mag with \\gsim 50\\%  completeness. It comprises a mixture of stellar generations, with 60\\% of the stars formed \\lsim\\,5\\,Gyr ago \\citep{cignoni11}. The young populations in NGC\\,346 consist mainly of low-mass pre--main-sequence (PMS) stars, identified from their positions on the color-magnitude diagram  \\citep[see][for a review on low-mass PMS stars in the Magellanic Clouds] {gouliermis12}. It is decidedly troublesome to determine an age for a cluster based on its PMS stars \\citep[][]{jeffries12, preibisch12}. Nevertheless, an isochronal age of $\\sim$\\,3\\,Myr has been established for NGC\\,346 by \\cite{sabbi08} by comparison with PMS evolutionary models. We complete the sample of low-mass PMS stars with the upper--main-sequence (UMS) stars ($m_{\\rm 555}-m_{\\rm 814} \\leq$\\,0.0\\,mag; 12\\lsim\\,$m_{\\rm 555}$\\lsim\\,17\\,mag; ages \\lsim\\,10\\,Myr), compiling a total sample of 5,150 stars. The UMS stars correspond to about 7\\% of the sample. The  map of our stellar inventory  is shown in Figure\\,\\ref{f:map}.  %%%%%%%%%  FIGURE  %%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{map-arcsec-reduced+fields.eps} \\caption{Map of the stars selected for our analysis, i.e., low-mass PMS and more massive upper--main-sequence (UMS) stars, as found in both F555W and F814W filters with {\\sc dolphot} photometry \\citep{dolphin00} based on ACS/WFC imaging of three pointings on the field of NGC\\,346. Red dots correspond to the UMS stars. The map covers the whole observed area. North is up, and east is to the left. Coordinates are given in seconds of arc in respect to a reference point, and are determined based on the drizzled ACS image in the F814W filter.} \\label{f:map} \\end{figure} %%%%%%%%%  FIGURE  %%%%%%%%%%%%%%%%%%%%%%%%%%  In a previous study \\citep[][from here on Paper\\,I]{schmeja09} we applied a cluster analysis based on the nearest-neighbor density method and we identified ten individual PMS stellar clusters in the region. We established that NGC\\,346 is a multi-clustered environment. This work also provided evidence of hierarchy in the PMS stellar clustering from a graph theory study with the minimum spanning tree and the analysis of the $\\cal{Q}$ parameter \\citep{cw04}. %Evidence of centrally concentrated stars on the main-sequence down to %$m_{\\rm 555} \\sim$\\,22\\,mag was also found. While these methods provided a unprecedented insight of the NGC\\,346 clustering, they were not able to quantitatively describe its complexity, or to characterize its self-similar behavior. In addition, the accuracy of the $\\cal{Q}$ parameter in interpreting a fractal structure has been challenged on the basis of the effect of projection in elongated star clusters (see \\citealt{bastian09} and \\citealt{cw09} for different accessions and viewpoints to the problem). The $\\cal{Q}$ parameter for the whole complex of NGC\\,346 was found equal to about 0.8, and thus cannot be used to conclude about the nature of stellar clustering in this region.  %In addition, the validity of the $\\cal{Q}$ parameter in establishing a structure %as fractal has been challenged on the basis of the effect of projection in elongated %star clusters (see \\citealt{bastian09} and \\citealt{cw09} for different accessions and %viewpoints to the problem).  We revisit the question of stellar clustering in NGC\\,346 with the application of a thorough cluster analysis. We first assess the topology of young stellar clustering in the region with the kernel density estimation technique, and the distribution function of stellar separations. We then decipher the clustering behavior of young stars in NGC\\,346 with the construction of their observed two-point correlation function, i.e., their autocorrelation function and its comparison to those from a series of simulated stellar distributions. We explore, thus, the limitations of this method and provide an accurate interpretation of the observed autocorrelation function. We, thus, carefully characterize the complex clustering behavior of young stars in NGC\\,346.  %We revisit the question of stellar clustering in NGC\\,346 with the application of a cluster analysis %based on the two-point correlation function, i.e., the autocorrelation function. We also assess the %topology of stellar clustering in the region with the application of the kernel density estimation %method. We compare the observed autocorrelation function to those from a series of simulated %stellar distributions in order to understand the limitations of this method and explore possible %interpretations of the observed stellar clustering. We characterize, thus, accurately the complex\u00ca %clustering behavior of young stars in NGC\\,346.  The paper is organized as follows. In Section\\,\\ref{s:tpcf} we present the stellar surface density map of NGC\\,346, and the distribution function of stellar separations. In Section\\,\\ref{s:acf} we present the autocorrelation function and discuss the first observable evidence that NGC\\,346 contains (at least) two components with very distinct distributions. We compare the autocorrelation function of NGC\\,346 with simulations of centrally condensed and self-similar stellar distributions in Section\\,\\ref{s:acfint}. We show that indeed the stellar distribution in the region is the result of two individual stellar components, a central condensed and an extended fractal distribution. In the same section we also constrain the basic parameters of these components using  dedicated simulations of mixed stellar distributions. Finally, in Section\\,\\ref{s:sum} we summarize our results and discuss their implications to our understanding of clustered star formation. Concluding remarks of our study are given in Section\\,\\ref{s:concl}. Additionally, in Appendix\\,\\ref{s:appendix} we present our library of simulated autocorrelation functions to be used for the interpretation of that observed in NGC\\,346, and in Appendix\\,\\ref{d3toeta} we provide an empirical calibration between three-dimensional and two-dimensional fractal dimensions over a wide range of  values.  %\\part{Spatial Distribution of Young Stars in NGC\\,346}  %--------------------------------------------------------------------------- ",
        "conclusions": "\\label{s:concl}  %The conclusions of our study on the clustering behavior of young stars in NGC\\,346 %can be summarized to the following points:  Concluding remarks, derived from our analysis, are summarized to the following:  \\begin{description}  %\\begin{itemize}  \\item Large star-forming complexes are excellent stellar `ecosystems' for the investigation of clustered star formation, because the represent the typical 100-pc scale of giant molecular clouds.  \\item The combined application of various cluster analysis tools is necessary for thoroughly characterizing young stellar clustering. In particular the autocorrelation function emerges as a  robust method, because, supported by the appropriate simulations, it is capable to distinguish different clustering styles.  \\item We determine that the clustering of young stellar populations in the star-forming complex NGC\\,346 has two distinct components; An extended self-similar, i.e., hierarchical, stellar distribution, and at least a centrally condensed young stellar cluster.  \\item Dedicated simulations of combined stellar distributions show that the condensed stellar component of fits to a spherical cluster, including 40\\% of the stars, with a core radius between 8\\arcsec and 9\\arcsec ($\\sim$\\,2.5\\,pc) and a power-law surface density profile with slope between 2.20 and 2.35. The remaining population is fractally distributed across the whole extent of the region with a fractal dimension $D_3$ between 2.20 and 2.36.  \\item Our findings suggest that the present clustering behavior of young stars in NGC\\,346 is the product of both star formation (turbulent-induced hierarchy), and early dynamical evolution (merging towards or dissolution of a centrally condensed cluster). Both processes seem to be still active after a time-scale of $\\sim$\\,3\\,Myr, i.e., the assumed evolutionary age of the region.  \\item The origin of this bimodal clustering behavior is not clearly understood. Considering that it influences the future evolution of the region, we assess that if the complete system is under disruption then possibly the central cluster was formed condensed. If NGC\\,346 contracts under its own gravity then this would imply that the central cluster may be the product of this contraction. We determine that the velocity dispersion limit for stability of the region as a whole is $\\sigma \\gsim$\\,6\\,-\\,11\\,km\\,s$^{-1}$. However, based on the heterogeneous clustering of the region, the {\\sl local} velocity dispersion of individual sub-structures should demonstrate variations.  %\\end{itemize}  \\end{description}  Considering the excellent coincidence of the fractal dimension derived for the stellar clustering with that measured and theoretically predicted for the turbulent interstellar medium, a natural step forward for this study is to understand how the stellar assembling process relates to the structuring behavior of the natal interstellar medium. This topic will be addressed in a subsequent study, where we compare the interstellar gas structure to that presented here for the young stars in NGC\\,346."
    },
    "1402/1402.0750_arXiv.txt": {
        "abstract": "The constraint imposed by magnetic helicity conservation on the alpha effect is considered for both magnetically and flow dominated self-organizing plasmas. Direct numerical simulations are presented for a dominant contribution to the alpha effect, which can be cast in the functional form of a total divergence of an averaged helicity flux, called the helicity-flux-driven alpha (H$\\alpha$) effect. Direct numerical simulations of the H$\\alpha$ effect are presented for two examples---the magnetically dominated toroidal plasma unstable to tearing modes, and the flow-dominated accretion disk. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5343_arXiv.txt": {
        "abstract": "% {} {% We present the first model that couples the formation of the corona of a solar active region to a model of the emergence of a sunspot pair. This allows us to study when, where, and why active region loops form, and how they evolve. } {% We use a 3D radiation MHD simulation of the emergence of an active region through the upper convection zone and the photosphere as a lower boundary for a 3D MHD coronal model. The latter accounts for the braiding of the magnetic fieldlines, which induces currents in the corona heating up the plasma. We synthesize the coronal emission for a direct comparison to observations. Starting with a basically field-free atmosphere we follow the filling of the corona with magnetic field and plasma. } {% Numerous individually identifiable hot coronal loops form, and reach temperatures well above 1 MK with densities comparable  to observations. The footpoints of these loops are found where small patches of magnetic flux concentrations move into the sunspots. The loop formation is triggered by an increase of upwards-directed Poynting flux at their footpoints in the photosphere. In the synthesized EUV emission these loops develop within a few minutes. The first EUV loop appears as a thin tube, then rises and expands significantly in the horizontal direction. Later, the spatially inhomogeneous heat input leads to a fragmented system of multiple loops or strands in a growing envelope. } {} ",
        "introduction": "}  In the process of emergence of magnetic field through the surface of the Sun a group of sunspots can form and the upper atmosphere is heated and filled with plasma. This forms  coronal loops prominently visible in extreme ultraviolet (EUV) and X-ray observations. In the early phase of theoretical investigations of  coronal loops, emphasis was put on their properties in an equilibrium state, i.e. once they have fully developed. For example, this led to the widely used scaling laws of \\cite{Rosner+al:1978} relating energy input and loop length to temperature and pressure of a static loop. While now time-dependent models including plasma flows are widely used, they mostly concentrate on the loop evolution in a magnetic field that changes only slightly or not at all \\citep[see e.g.\\ review by][]{Reale:2010}. This applies to 1D models for the dynamic evolution of loops, e.g.\\ accounting for siphon flows \\citep{Boris+Mariska:1982}, in response to intermittent nanoflare heating \\citep{Hansteen:1993}, or concerning catastrophic cooling \\citep{Muller+al:2003}. Also 3D models of the coronal loop structure considered either a constant or slowly changing magnetic field in the photosphere \\citep{Gudiksen+Nordlund:2002, Bingert+Peter:2011,Lionello+al:2013}. In all these models, the loops form and evolve in a preexisting coronal magnetic field above a well developed active region. However, during the early formation stage of an active region, when the first coronal loops form,  flux emergence is still significant. Using a magneto-frictional approach, \\cite{Cheung+DeRosa:2012} followed the whole evolution of an active region. In particular, they determined the coronal magnetic field  in response to the flux emergence and used proxies based on  currents to obtain a rough estimate of the coronal emission. The goal of our study is to go one step beyond and follow the formation of the first loops in an emerging active region in a self-consistent model, properly accounting for the plasma properties and the magnetic field. This requires the treatment of an energy equation including radiative losses, heat conduction, and plasma heating. The synthesized coronal emission can be directly compared to observations.    Over the last decade 3D models showed that in the magnetically closed {\\em active region} corona the heat input most probably is related to the tangling of the magnetic field in the photosphere, either through fieldline braiding by the horizontal motions of granulation \\citep{Parker:1972, Parker:1983} or flux-tube tectonics \\citep{Priest+al:2002}. In both cases currents are induced in the corona and are subsequently dissipated.  In a numerical experiment \\citet{Galsgaard+Nordlund:1996} could show that the driving of the magnetic field from the photospheric boundaries can indeed lead to significant build-up of currents to heat the corona through Ohmic dissipation. In a more realistic setup \\citet{Gudiksen+Nordlund:2002,Gudiksen+Nordlund:2005a,Gudiksen+Nordlund:2005b} showed that a million K hot loop-dominated corona forms above an active region. They used an observed magnetogram and the fieldline braiding was driven by artificial granulation-like motions at their bottom boundary. The induced currents $j$ are then converted to heat through Ohmic dissipation $H_{\\rm{Ohm}}\\propto\\eta{j^2}$, where $\\eta$ has a functional dependence on $j$. Transition region EUV emission lines synthesized from these models match major observed features such as the transition region redshifts  \\citep{Peter+al:2004,Peter+al:2006} or temporal variability \\citep{Peter:2007}. Using a constant value for $\\eta$, \\citet{Bingert+Peter:2011} studied the temporal variability of the heat input in detail. Such models show that the energy is predominantly deposited in quantities of about $10^{17}$ J \\citep{Bingert+Peter:2013}, consistent with the predictions of nanoflares by \\cite{Parker:1988}. Most recently, observed coronal structures could be reproduced by the 3D coronal models based on the observed magnetic field and flows in the photosphere alone \\citep{Bourdin+al:2013}, and the constant cross section of loops was reproduced \\citep{Peter+Bingert:2012}. While the above models are all based on DC-type Ohmic heating, recently \\citet{Ballegooijen+al:2011} studied the AC energy input by Alfv\\'en wave turbulence in a coronal magnetic flux tube using a reduced MHD approach and found it sufficient to heat the hot coronal loops.  On smaller spatial scales, more representative for network or plage patches, \\citet{Martinez+al:2008, Martinez+al:2009} studied the process of a twisted flux tube emerging from the convective layer into the lower corona. In these realistic models the radiative transfer in the lower atmosphere is treated as well as the heat conduction and radiation needed in the coronal part. Descriptions of the flux tube emergence into the corona on larger scales, more representative of an active region, had to omit the realistic description allowing for the synthesis of photospheric and coronal emission \\citep[e.g.][]{Fan:2001, Abbett+Fisher:2003, Manchester+al:2004, Archontis:2004,Magara:2006}. \\citet{Hurlburt+al:2002} used a 2D axisymmetric magneto-convection sunspot model to construct the coronal magnetic field above. In a collection of 1D static models along the coronal fieldlines they then derived the EUV and soft X-ray emission and compared it to observations. This model is still quite idealized in the sense that the corona is described by a series of 1D equilibrium models along  magnetic fieldlines, and the heat input is determined by a (freely chosen) fraction of the Poynting flux that is generated by the magneto-convection.  In recent years it became possible to study the fine structure of sunspots \\citep{Heinemann+al:2007} and simulate a whole sunspot and its surrounding granulation in 3D models from the convection zone to the photosphere \\citep{Rempel+al:2009.sci}. This paved the road to investigate the emergence of a flux tube emerging through the surface resulting in a sunspot pair that is at the heart of an active region \\citep{Cheung+al:2010}. These models are realistic in the sense that they include all the relevant physics, in particular the radiative transfer, in order to produce synthetic images from the simulation that closely resemble actual observations. A variant of this latter model (see \\sect{S:model_flux}) will provide the lower boundary for our model of the corona above an active region.  The goal of our study is to use a realistic model of the emergence of a sunspot pair in the photosphere as a lower boundary to drive a coronal model of the active region corona above. A single model encompassing a whole active region from the interior to the upper atmosphere is not yet possible. Solving both the radiative transfer in the lower atmosphere and  the anisotropic heat conduction in the corona in a computational domain covering a full active region is too demanding in terms of computing resources. Instead, we couple our coronal model to the flux-emergence simulation by driving the coronal model at the bottom boundary with the results from the flux-emergence simulation in the (middle) photosphere.  The model setup is described in \\sect{S:model}, before we present the appearance of the model corona in \\sect{S:appear}. A detailed analysis along the coronal loop that forms first during the flux emergence is given in \\sect{S:loop} followed by a discussion of the  three-dimensional nature of the loop in \\sect{S:sect}. Finally we discuss the trigger of the coronal loop formation in \\sect{S:trigger}. ",
        "conclusions": "In this paper, we presented a coronal model of an emerging active region driven by a simulation of magnetic flux emergence from the convection zone through the photosphere. The magnetic field  expands  into the corona, while a pair of simulated spots forms in the photosphere. Ohmic dissipation heats the coronal plasma, while heat conduction along magnetic fieldlines, radiative losses through optically thin radiation, and flows  carry away the energy input. The  treatment of the full energy balance ensures that the coronal pressure is set self-consistently and allows us to synthesize the  EUV emission from the model corona.  Once sufficient magnetic flux was emerged through the surface and the coalescence of small-scale magnetic patches formed large-scale magnetic patches turning into sunspots, the first EUV coronal loops form within minutes. The EUV  loop rises upwards, expands significantly in the horizontal direction, and, most importantly, fragments into several individual EUV structures, i.e. the changing heat input produces new strands in a growing envelope.  The energy input is driven by the advection of the magnetic field in the photosphere, i.e. by the horizontal convective motions. Connected by magnetic fieldlines through the corona, the regions of enhanced Poynting flux at one end induce an increase of the Poynting flux at the other end at the coronal base. The upward directed Poynting flux leads to an increased energy input giving rise to the heating of the coronal plasma and the enhancement of the pressure due to the evaporative upflows.  The emerging magnetic field hosting the forming loop rises into the ambient magnetic field and currents also build up near the upper part. These contribute to the Ohmic heating   in the top part of the loop leading to a nearly constant heat input along the loop.  In its early evolution  the coronal loop behaves (at least concerning the energy budget) similarly as a conventional 1D loop model would predict if we would prescribe the energy input. However, in the later stages the loop shows its true 3D nature. The horizontal magnetic expansion and in particular the fragmentation of the EUV emission are a clear indication that a 1D model would not be sufficient to describe a newly forming emerging loop. In  the cross-sectional cut perpendicular to the EUV loop, the temperature and the density structure are comparably smooth but not exactly cospatial. This gives rise to the fragmented appearance of the loop in EUV emission with threads (or loop-fragments, or strands) with diameters much smaller than the typical spatial structures in temperature or density.  Our model of the formation and evolution of a EUV coronal loop in an emerging active region sheds new light on our understanding of coronal loop formation. A further analysis of this and more advanced numerical experiments will have to investigate the differences (and similarities) of the evolution of coronal loops seen in different wavelength bands, in particular towards X-rays, and how the forming loops would appear in spectroscopic observations. Of particular interest will be the further investigation of the evolution of the magnetic field structure in relation to the spatial structure of the synthesized coronal emission."
    },
    "1402/1402.0085_arXiv.txt": {
        "abstract": "We analyze a large set of new and archival photometric and spectroscopic observations of LMC X-3 to arrive at a self-consistent dynamical model for the system.  Using echelle spectra obtained with the MIKE instrument on the 6.5m Magellan Clay telescope and the UVES instrument on the second 8.2m Very Large Telescope we find a velocity semiamplitude for the secondary star of $K_2=241.1\\pm 6.2$ km s$^{-1}$, where the uncertainty includes an estimate of the systematic error caused by X-ray heating. Using the spectra, we also find a projected rotational velocity of $V_{\\rm rot}\\sin i=118.5\\pm 6.6$ km s$^{-1}$.  From an analysis of archival $B$ and $V$ light curves as well as new $B$ and $V$ light curves from the SMARTS 1.3m telescope, we find an inclination of $i=69.84\\pm 0.37^{\\circ}$ for models that do not include X-ray heating and an inclination of $i=69.24\\pm 0.72^{\\circ}$ for models that incorporate X-ray heating. Adopting the latter inclination measurement, we find masses of $3.63\\pm 0.57\\,M_{\\odot}$ and $6.98\\pm 0.56\\,M_{\\odot}$ for the companion star and the black hole, respectively. We briefly compare our results with earlier work and discuss some of their implications. ",
        "introduction": "LMC X-3 was the second black hole, after Cygnus X-1, to be established via dynamical observations (Cowley et al.\\ 1983).  The large mass function of the host binary, $f(M) = 2.3 \\pm 0.3~M_{\\odot}$, the absence of X-ray eclipses, and the estimated mass of the B3 V secondary ($4-8~M_{\\odot}$) allowed Cowley et al.\\ to conclude that ``the most probable mass'' of the black hole is $M \\sim 9 M_{\\odot}$, with a plausible lower limit of $M > 7~M_{\\odot}$.  A subsequent study of the 1.7-day optical light curve constrained the inclination $i$ to lie in the range $\\sim64^{\\circ}$ to $70^{\\circ}$, and it also provided a firm lower limit on the mass of the primary, $M > 3.5~M_{\\odot}$ (Kuiper et al.\\ 1988), thereby confirming that it is indeed a black hole (e.g., Kalogera \\& Baym 1996).  Recently, in the context of a far-UV study of LMC X-3 (Song et al.\\ 2010), and using much of the data discussed herein, we reported a precise ephemeris and radial velocity amplitude for the secondary.  It is a challenge to establish a reliable dynamical model of the LMC X-3 system because both the optical counterpart and the X-ray source are highly variable.  In addition to the $\\approx 0.15$ mag ellipsoidal variations (Kuiper et al.\\ 1988), the optical flux varies chaotically by up to $\\approx 1$ mag (Brocksopp et al.\\ 2001), while the {\\it RXTE} PCA shows that the X-ray luminosity ranges over a factor of at least 2,500 (Smale \\& Boyd 2012).  Furthermore, the optical variability (which is most relevant for a dynamical study) is not simply slaved to the X-ray variability as it is in most transient black hole binaries: In a typical transient system, the optical emission is generated directly and promptly by X-ray reprocessing in the outer accretion disk or in the X-ray-heated face of the secondary (van Paradijs \\& McClintock 1995).  While both of these effects do occur in LMC X-3, the extra complicating factor for this source is the variable mass accretion rate in the outer disk that produces an additional, large-amplitude component of optical and (delayed) X-ray variability (Brocksopp et al.\\ 2001; Steiner et al. 2013).   The absorbing column depth to LMC X-3 is stable and extraordinarily small, $N_{\\rm H} = 3.8_{-0.7}^{+0.8} \\times 10^{20}$~cm$^{-2}$ (Page et al.\\ 2003) and, correspondingly, the optical extinction is low, $A_{\\rm V} \\approx 0.2$~mag.  Also, the X-ray spectrum is normally disk-dominated (Wilms et al.\\ 2001; Steiner et al.\\ 2010).  This combination of low absorption and a soft spectrum, combined with the persistence of the source, has made LMC X-3 a touchstone for testing accretion disk models (Davis et al.\\ 2006; Kubota et al.\\ 2010; Straub et al.\\ 2011), and for stringently testing the stability of the inner disk radius as a foundation for the measurement of black hole spin (Steiner et al.\\ 2010).  Davis et al.\\ (2006) have used the continuum-fitting method to estimate the dimensionless spin parameter of the black hole primary in LMC X-3 to be $a_* \\lesssim 0.3$, a value that depends strongly on the provisional estimates of the black hole mass $M$ and inclination $i$ that are discussed above (Cowley et al.\\ 1983; Kuiper et al.\\ 1988).  In this paper, which is based on an extensive collection of spectroscopic and photometric data, we present a new dynamical model for LMC X-3.  In a companion paper (Steiner et al. 2014), we make use of the precise values of $M$ and $i$ reported herein, the secure distance to LMC X-3, and extensive archival X-ray data (Steiner et al.\\ 2010) to obtain a firm estimate of the black hole's spin.  This paper is organized as follows.  We discuss the new and archival spectroscopic observations in \\S2 and the new and archival photometric observations in \\S3.  In \\S4 we discuss the analysis of the spectra including the measurement of the radial velocities, the rotational velocity, and the detection of emission lines in some of the spectra. In \\S5 we first present our extensive collection of light curve data, and we then present our dynamical model.  We discuss several topics and offer our conclusions in \\S6. ",
        "conclusions": "Our mass for the black hole of $6.98\\pm 0.56\\,M_{\\odot}$ is considerably more precise than earlier estimates for three principal reasons: First, we have far more radial velocity data of high quality for the secondary star, which yield a much improved determination of its $K$-velocity. Second, we have obtained the first accurate measurement of the projected rotational velocity of the secondary, which is an important constraint on the dynamical model. Third, we have analyzed a much larger body of ellipsoidal light curve data than any previous study (e.g., Kuiper et al. 1988).  In their pioneering work on LMC X-3, Cowley et al. (1983) concluded that the most plausible mass of the black hole was $\\sim 9\\,M_{\\odot}$ with a lower limit of $M>7\\,M_{\\odot}$.  To reach this conclusion, Cowley et al.\\ had to assume a mass for the secondary star.  Based on their spectral classification of the star as B3V, they deduced its mass to be $4-8\\,M_{\\odot}$, greater than our value of $M_2=3.63\\pm 0.57\\,M_{\\odot}$.  Estimating the mass based on the spectral type is problematic for a star that has lost an indeterminate amount of mass. On a related note, our measurement of the surface gravity of the star, $\\log g=3.740 \\pm 0.020$, shows that it matches more closely the gravity of a giant (class III) than the gravity of a main sequence star (class V).  For example, the nominal masses and radii for a B3V/B5V star given by Ostlie \\& Carroll (1996) correspond to surface gravities of $\\log g=4.07/3.98$, while their surface gravity for a B5III star is $\\log g=3.68$.  Soria et al.\\ (2001) and Val-Baker et al.\\ (2007) provided mass estimates after making corrections to the $K$-velocity of the secondary star to account for X-ray heating.  Based on the system colors as measured by the Optical Monitor on {\\em XMM-Newton}, Soria et al.\\ concluded that the secondary star is a B5 subgiant that fills its Roche lobe.  Using an evolutionary model they determined a mass of $M_2=4.7\\,M_{\\odot}$ for the secondary star.  Next, Soria et al.\\ applied a $K$-correction of $\\Delta K=-30\\pm 5$ km s$^{-1}$ to the $K$-velocity from Cowley et al.\\ (1983) to arrive at the ``true'' $K$-velocity of $205\\pm 12$ km s$^{-1}$. The sign of their $K$-correction is negative because Soria et al. used the procedure that assumes heating by ``soft'' X-ray photons (Section 5.2.2). Finally, Soria et al.\\ used their secondary star mass and corrected $K$-velocity to arrive at their lower limit of $M > 5.8\\pm 0.6\\,M_{\\odot}$.   Val-Baker et al.\\ (2007) measured a $K$-velocity of $K_2=242.4\\pm 4.3$ km s$^{-1}$ from the five UVES spectra. Their $K$-correction of 14.3 km s$^{-1}$ is positive because they used the procedure that assumes irradiation by ``hard'' photons; their adopted $K$-velocity is therefore $K_2=256.7\\pm 4.9$ km s$^{-1}$ . By comparing their spectra to standard star spectra, Val-Baker et al.\\ determined a spectral type of B3V at phase 0.22 when X-ray heating is significant, and B5V at phase 0.86, when X-ray heating is slight.  Converting these spectral types into temperatures using Kurucz models, they arrived at effective temperatures of 15,400 K for the unheated face and 18,700 K for the heated face.  This temperature change was used to then find their $K$-correction.  Finally, after adopting a nominal mass of $5.9\\,M_{\\odot}$ for a B5V star (e.g.\\ Ostlie \\& Carroll 1996), Val-Baker found a mass range for the black hole of $9.5\\le M \\le 13.6\\,M_{\\odot}$.  Soria et al.\\ (2001) and Val-Baker et al.\\ (2007) not only disagree by a factor of two on the magnitude of the $K$-correction, they even disagree on its sign.  As noted earlier, the disagreement on the sign of the effect results from Soria et al.\\  assuming that the X-ray spectrum illuminating the star is soft, while Val-Baker et al.\\ assumed it is hard. Setting aside the difference in the signs for the moment, we believe that the magnitudes of these corrections are too large [where $|\\Delta K|=30$ km s$^{-1}$ for Soria et al.\\ (2001) and $|\\Delta K|=14.3$ km s$^{-1}$ for Val-Baker et al.\\ (2007)]. In contrast, consideration of our much more extensive data set shows that X-ray heating has a modest effect on the $K$-velocity (Section 4.1), which is probably because the star is shielded by the outer part of the accretion disk and because the X-ray emitting region has a disk-like geometry. Val-Baker et al.\\ inferred a temperature difference of 3300 K between the heated and unheated faces of the star. However, as shown in Figure \\ref{temperaturetrend}, we measure a temperature range of less than half their value.  The result of Val-Baker et al.\\ is subject to the complication that the X-ray luminosity of the system (as measured by the {\\em RXTE} ASM) varied by a factor of two over the course of their UVES observations (see Table \\ref{BAT}). Because the $K$-correction is strongly supralinear with temperature, our correction for 1500 K is much smaller than that of Val-Baker et al. for 3300K.    The black hole mass of $6.98\\pm 0.56\\,M_{\\odot}$ is entirely consistent with the masses of black hole in transient systems with low-mass secondaries ($7.8 \\pm 1.2~M_{\\odot}$; \\\"{O}zel et al. 2010; Farr et al. 2011). In contrast, the black hole masses in the persistently-bright, wind-fed systems with high mass secondaries (e.g.\\ M33 X-7, LMC X-1, and Cyg X-1) are considerably larger (Orosz et al.\\ 2007, 2009, 2011).  Considering the mass of the black hole and the following three other properties of LMC X-3, we classify the system as a member of the transient black hole binaries that are fed by Roche-lobe overflow, rather than as a member of the persistent systems whose black holes are wind-fed by massive O-type or Wolf-Rayet secondaries (McClintock et al. 2013): (1) Mass transfer via Roche-lobe overflow is expected given the mass ratio reported herein ($M/M_2 \\approx 2$). (2) Although normally X-ray bright, LMC X-3 is highly variable and on occasion enters a prolonged low-intensity state that is dominated by a hard power-law component (Wilms et al. 2001; Smale \\& Boyd 2012).  (3) Likewise, the low spin of the black hole (Davis et al. 2006; Steiner et al. 2014) distinguishes LMC X-3 from the persistent black holes and is consistent with that of other transient black holes (McClintock et al. 2013).   One obvious difference between LMC X-3 and a typical transient system like A0620-00 is that LMC X-3 has never been reported to be in a true X-ray quiescent state with $L_{\\rm x}\\lesssim 10^{33}$ ergs s$^{-1}$.  The mass transfer rate in LMC X-3 is high compared to a typical transient because the relatively massive secondary is presumably evolving on a nuclear timescale.  The higher mass transfer rate results in much shorter recurrence timescales, according to the disk instability model (Cannizzo, Chen, \\& Livio 1995; Lasota 2001).  In the context of this model, perhaps the recurrence time is so short that LMC X-3 is perpetually in an outburst state.  Finally, although a detailed evolutionary model is beyond the scope of this paper, we briefly speculate on the future of LMC X-3.  Since the present-day mass of the secondary star is less than that of the black hole, the mass transfer is expected to be both thermally and dynamically stable (Tauris \\& van den Heuvel 2006).  Mass transfer from the lower mass object to the higher mass object results in an expansion of the orbit (e.g., the period and the separation both increase).  The secondary star's radius would also increase, as it continues to fill its Roche lobe.  The mass that is transferred through the accretion disk will tend to increase the spin of the black hole.  By the time core H-burning ceases, the star may have transferred a considerable amount of its current mass to the black hole.  Thus, it appears that LMC X-3, will end up with a relatively massive black hole, a long orbital period, appreciable spin, and a low mass companion.  This future version of LMC X-3 will likely bridge the gap between most transient systems and the extreme system GRS 1915+105 (McClintock et al.\\ 2006; Steeghs et al.\\ 2013)."
    },
    "1402/1402.3938_arXiv.txt": {
        "abstract": "The UV properties of old stellar populations have been subject of intense scrutiny from the late sixties, when the UV-upturn in early type galaxies was first discovered.  Because of their proximity and relative simplicity, Galactic globular clusters (GGCs) are ideal local templates to understand how the integrated UV light is driven by hot stellar populations, primarily horizontal branch stars and their progeny. Our understanding of such stars is still plagued by theoretical uncertainties, which are partly due to the absence of an accurate, comprehensive, statistically representative homogeneous data-set. To move a step forward on this subject, we have combined the {\\it HST} and {\\it GALEX} capabilities and collected the largest data-base ever obtained for GGCs in UV. This data-base is best suited to provide insights on the HB second parameter problem and on the first stages of GCs formation and chemical evolution and to understand how they are linked to the observed properties of extragalactic systems. ",
        "introduction": "It is fair to say that the last frontier of our growing understanding of the physics of old stellar populations resides in the ultraviolet (UV). The behavior of old stellar populations in the UV has puzzled astronomers for almost four decades now, and in spite of major recent progress, there are still important gaps in our understanding of the nature of stars that dominate the integrated light of old stellar populations in the UV, particularly the far-UV (FUV, Ferraro et al. 1998; O' Connell 1999, Moehler 2001; Catelan 2009; Dalessandro et al. 2011, 2012, 2013a; Schiavon et al. 2012). \\\\ In distant extragalactic systems one can ordinarily observe only the integrated light of unresolved stellar populations, from which the hope is to gain knowledge about the underlying stellar population. Galactic globular clusters (GGCs) play an important role in understanding the integrated UV colors of extragalactic systems, especially the so called \"UV-upturn\" observed in the spectral energy distributions of elliptical galaxies (Code \\& Welch 1979; de Boer 1982; Bertola et al. 1982; Greggio \\& Renzini 1990; O'Connell 1999). First of all, GCs are the closest example in nature to a single stellar population: a system of coeval stars with similar chemical composition.  Moreover GGCs span a large range of metallicities, a small range of ages, and perhaps some range of helium abundance. Hence they can be used to test the stellar evolution theory, which in turn is one of the basic ingredients of the models used to interpret the integrated light of distant galaxies. GGCs are relatively nearby objects (more than $\\sim90\\%$ are located at distances $r<30$\\,kpc), so their populations can be easily resolved.  With typically more than 100,000 stars, even relatively short-lived evolutionary stages are sampled. \\\\ In general the main contributors to the UV emission from any stellar system are the hottest stars. Indeed, blue horizontal branch (HB) stars and their progeny are well known to be among the hottest stellar populations in GCs and contribute substantially to the UV radiation observed from old stellar systems (Welch \\& Code 1972). The hottest HB stars (extreme HB, EHB) have such a small envelope mass that most of their post-He-core burning phase takes place at high effective temperature ($T_{\\rm eff}$), during the so called \"AGB-manqu\\'e phase\", and these stars never return to the asymptotic giant branch (AGB). Another group of UV-bright stars is that of post-early AGB stars, which after a brief return to the AGB, spend the bulk of their helium shell burning phase at high $T_{\\rm eff}$. In systems with only red HB a small floor level of $FUV$ is provided by post-AGB stars, which evolve to the AGB phase with an higher envelope mass where they undergo thermal pulses and eventually lose their envelopes moving at higher temperatures at constant luminosity. \\\\ The relative contributions of the various types of stars and the factors that might lead to larger or smaller populations of UV-bright stars have remained an open question (Greggio and Renzini 1990; Dorman et al. 1995; Lee et al.  2002; Rich et al. 2005; Sohn et al. 2006). The capability to predict the relative contribution of various UV emitters in GCs is strongly linked to our knowledge of the physical mechanisms shaping the HB morphology. The scientific community agrees from nearly 50 years about the fact that the principal parameter governing the shape of HBs in GCs is metallicity. The general rule is that metal-rich systems have red HBs, while in the metal-poor ones stars are distributed on average at higher effective temperatures. However several exceptions to this general trend are observed: remarkable cases are those of NGC~6388 and NGC~6441 (Rich et al. 1997), which despite their metallicity ([Fe/H]$\\sim-0.6$) show some of the bluest HBs known in GGCs (Busso et al. 2007; Dalessandro et al. 2008). Moreover several clusters, sharing similar metal content, reveal different HB morphologies, typical cases being the pairs NGC~5927 - NGC6388 at high metallicities ([Fe/H]$\\sim-0.4$), M~3 - M~13 at intermediate-metallicity regime ([Fe/H]$\\sim-1.5$) and M~15 - M~92 at low-metallicities ([Fe/H]$\\sim-2.3$). These noticeable exceptions have required the introduction of a second (Freeman \\& Norris 1981) and possibly a third parameter in order to explain the HB distributions in all GGCs. What we can call now the \\emph{i-th parameter problem} is still a hot topic, as stressed by several authors, we recall the reader to Catelan 2009 for a nice review (see also Dotter et al. 2010 and Gratton et al. 2010). Despite the huge efforts made to address this problem, its solution is not obvious and still different scenarios are proposed. One of the reasons that complicates the identification of the mechanisms -- other than metallicity -- at work in shaping the observed luminosity and effective temperature distribution of stars along the HB is that there are many possible culprits (mass-loss, age, helium abundance ...; see Rood 1973 for example) and some of them are not well constrained from theory. Age has been identified as the natural global second parameter by many authors in the past years (Lee et al. 1988, 1990; Lee, Demarque \\& Zinn 1994; Sarajedini \\& King 1989; Dotter et al. 2010; Gratton et al. 2010).  According to this interpretation older clusters tend to have bluer HBs, while younger ones should have on average redder HB morphologies. \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.4]{fig7_pr.ps} \\caption{Optical CMD of the HB of NGC~2808 compared to theoretical ZAHBs for $Y=0.248, 0.30$ and 0.40. The dashed lines show the effect of neglecting radiative levitation in the ZAHB bolometric corrections when $T_{\\rm eff}$ is larger than 12000 K. } % \\end{center} \\end{figure} This scenario appeared in agreement with the picture for the Galaxy formation and its early evolution (Searle \\& Zinn 1978; Zinn 1985). Still, age is not able to explain exhaustively the HB morphology. Detailed cluster to cluster comparisons have shown that there are systems with similar iron content and age, but remarkably different HB morphologies. The necessity of at least a third parameter transpires also from Dotter et al. (2010) and Gratton et al. (2010) analyses, in the form of either the luminosity cluster density or stellar density ($\\log(\\rho)$) -- as already suggested by Fusi Pecci et al. (1993) -- which might correlate with the hot extension of the HBs, or a variation of the initial helium abundance (Y), respectively.\\\\ GGCs are therefore crucial local templates for comparison with integrated properties of distant extragalactic systems. Comparing features in the color-magnitude diagrams (CMDs) of well known and resolved GGCs with integrated quantities can lend important {\\it model independent} insights into the nature of extragalactic systems. ",
        "conclusions": "Understanding the origin and the frequency of hot stars is not simply a problem of understanding the evolution of old, low mass stars, but it has important implications on the interpretation of the integrated properties of galaxies. Indeed hot stars have been suggested to be responsible of the UV-upturn in the spectra of elliptical galaxies and bulges (Greggio \\& Renzini 1990). \\\\ As a part of a large project aimed at studying the properties of hot stellar populations in GGCs as a tool to interpret the properties of unresolved stellar populations in other galaxies, we have presented results obtained with a large UV database built with two complementary instruments: HST and GALEX.   \\\\ We have shown that analyses based on a suitable combination of UV photometry and synthetic HB simulations can provide not only insights on the HB second parameter problem, but in general they can potentially give some clues on the first stages of GCs formation and chemical evolution (Dalessandro et al. 2011, 2013a). In fact while metallicity and age are believed to be the main general parameters shaping the HB morphology (Lee et al. 2002; Dotter et al. 2010; Gratton et al. 2010), there are cases that can be reproduced by only invoking different initial He abundances. In this sense a nice example is represented by the well known triplet M~3 - M~13 - M~79. In fact these GGCs share the same metallicity and age but have incredibly different HB morphologies. We find (Dalessandro et al. 2013a) differences of initial He abundances $\\Delta Y_{\\rm max}\\sim0.02-0.04$, in particular M~13 displays the largest value ($Y_{\\rm max}\\sim0.30$, M~3 the smallest one ($Y_{\\rm max}\\sim0.27$) and M~79 an intermediate case  ($Y_{\\rm max}\\sim0.27$). These values seem to qualitatively correlate with the differences in temperature (color) extensions of their HBs. Interestingly they correlate also with the observed range of light-element variations. \\\\ Since UV colors (FUV in particular) are primarily driven by the number of HB and post-HB stars and their temperature distribution, we should expect He to have a quite important role on the integrated properties of GCs. Indeed, from the comparison between GGCs and those belonging to the Sagittarius dwarf, M31 and M87, different behaviors emerge. In fact the clusters associated with the Sagittarius dwarf are on average redder than the MW ones, the M~31 clusters have colors which are comparable to those of the GGCs, while the M~87 star systems are bluer. We note that there may be a possible trend between the mass of the host galaxy and the color distribution of its globulars, in the sense that the higher is the galaxy mass, the bluer are the GC UV colors.  In fact M~87 (with the bluer systems) is a super-giant elliptical that is about two orders of magnitude more massive than the Milky Way ($1.7\\times10^{13}<M/M_{\\odot}<4.0\\times10^{13}$; Fabricant et al. 1980), while Sagittarius ($\\sim1.6\\times10^8 M_{\\odot}$; Law \\& Majewski 2010) with the reddest sample of GCs (although quite small), is a dwarf galaxy, and M31 ($3.7\\times10^{11}<M/M_{\\odot}<2.5\\times10^{12}$; C{\\^o}t{\\'e} et al.2000) and the Galaxy ($2.4\\times10^{11}<M/M_{\\odot}<1.2\\times10^{12}$; Little \\& Tremaine 1987; Kochanek 1996) representing intermediate cases. We argue that most of the observed differences between colors involving the $FUV$ band are explainable invoking different Helium contents.  This would lead us to speculatively think that galaxies with larger masses may have, on average, more He-rich populations. In that case, He abundance differences could be a by-product of chemical evolution differences, in some way connected to the mass of the host galaxy. This could be also connected with the formation and dynamical history of clusters in galaxies with different masses, as suggested by Valcarce \\& Catelan (2011; see also Carretta et al. 2014). In particular they argue that clusters hosted by more massive galaxies are more likely to undergo a more complex history of star formation thus having a larger spread in stellar populations properties."
    },
    "1402/1402.4804_arXiv.txt": {
        "abstract": "The radio--loud quasar SDSS J114657.79+403708.6 at a redshift $z$=5.0 is one of the most distant radio--loud objects. The IR--optical luminosity and spectrum suggest that its black hole has a very large mass: $M=(5\\pm 1) \\times 10^9 M_\\odot$. The radio--loudness (ratio of the radio to optical flux) of the source is large (around 100), suggesting that the source is viewed at small angles from the jet axis, and could be a blazar. The X--ray observations fully confirm this hypothesis, due to the high level and hardness of the flux. This makes SDSS J114657.79+403708.6 the third most distant blazar known, after Q0906+693 ($z=5.47$) and B2 1023+25 ($z=5.3$). Among those, SDSS J114657.79+403708.6 has the largest black hole mass, setting interesting constraints on the mass function of heavy ($>10^9 M_\\odot$) black holes at high redshifts. ",
        "introduction": "\\label{sec-intro}  Since the radiation produced by relativistic jets is strongly boosted along the jet direction, objects whose jet is pointing at us are very bright, and can be seen up to high redshifts. The black hole at the center of the powerhouse of these sources can be very massive, sometimes exceeding $M=10^{10}M_\\odot$ (Ghisellini et al. 2009; 2010a). % Therefore the hunt for high redshift blazars is an important field of research, allowing the census of heavy black holes in the early Universe. This can confirm or challenge existing theories of black hole formation, and even more so if we associate the presence of relativistic jets with a large black hole spin. In this case, in fact, the efficiency of accretion $\\eta$ (the fraction of accreted mass transformed into radiation) is higher than for a non--rotating black hole, reaching a value of $\\eta=0.3$ for maximally rotating {\\it accreting} black holes (see Thorne 1974). This implies that the Eddington luminosity is reached with a smaller accretion rate than for a non--rotating black hole. If the system is Eddington limited, this in turn implies a slower black hole growth. If a black hole seed starts to accrete at a redshift $z=20$ at the Eddington rate maintaining a large spin, it can reach a billion solar masses only at $z<4$ even for a seed mass as large as $10^6 M_\\odot$ (see e.g. Ghisellini et al. 2013). The very fact of the existence of radio--loud sources with $M>10^9 M_\\odot$ at $z>4$ is thus a problem.  These issues motivate our search of high redshift blazars, keeping in mind that for each blazar (i.e. viewing angle smaller than $1/\\Gamma$, where $\\Gamma$ is the jet bulk Lorentz factor) there must be other $2\\Gamma^2$ sources whose jets are pointing in other directions.  Up to now there are two blazars known at $z>5$: Q0906+6930 ($z=5.47$, Romani et al. 2004; Romani 2006), and B2 1023+25 ($z=5.3$, Sbarrato et al. 2012, 2013). The spectral energy distribution (SED) of these sources reveals both the thermal (i.e. strong optical emission lines and continuum) and the boosted non--thermal components. As in the majority of very powerful and high--$z$ blazars, the thermal disk emission becomes visible since it stands between the two non--thermal humps (the synchrotron one peaking in the sub-mm, and the high energy in the $\\sim$MeV band; Ghisellini et al. 2010a; 2010b). In these sources the  X--ray spectrum is hard [i.e. $\\alpha_X \\sim 0.5$, assuming $F(\\nu) \\propto \\nu^{-\\alpha_{x}}$], and this, together with a relatively strong X--ray to optical flux ratio, can be taken as a signature of the blazar nature of the source.  In this letter we suggest that SDSS J114657.79+403708.6, a radio--loud AGN at $z=5.005$, is a blazar, i.e. the viewing angle is smaller than $1/\\Gamma$. Evidences for its blazar nature include its relatively large radio--loudness and its very large X--ray luminosity and hard X--ray spectrum, as measured by a pointed {\\it Swift} observation. Infrared data collected by the WISE satellite (Wright et al. 2010), together with the Sloan Digital Sky Survey (SDSS; York et al. 2000) spectrum  allowed to constrain the properties of the thermal emission.  In this work, we adopt a flat cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": "\\label{sec-discussion}   In this letter we propose that the radio--loud, high redshift quasar SDSS 1146+403 is a blazar. If so, it is the third most distant blazar known up to now, with redshift $z=5.005$, that corresponds to a cosmic age 1.1 billion years. Despite this young age, the black hole of SDSS 1146+403 managed to grow to 5 billion solar masses.  Both its thermal and non--thermal components are very luminous. In agreement with the blazar sequence (Fossati et al. 1998) the two broad non--thermal humps peak at small frequency. In particular, the hard X--ray spectrum and the upper limit in the $\\gamma$--ray band constrain the high energy component to peak in the MeV region of the spectrum. Therefore this source, along with the similar other powerful blazars, should have a relatively large hard X--ray flux and would have been an ideal target for hard X--ray instruments such as the focussing hard X--ray telescope {\\it NuStar} (Harrison et al., 2013). The great sensitivity over its energy range [5--80 keV] would enable it to detect the hard X--ray spectrum of this source, even if it cannot directly observe the peak of the high energy hump.  \\vspace{0.3 cm} It is possible to roughly estimate the number of these objects and their spatial density.  The comoving density of heavy black holes at high redshifts of radio--loud sources has been studied by Volonteri et al. (2011), based on the 3 years BAT catalog and the blazar luminosity function, in hard X--rays, derived by Ajello et al. (2009), and modified (beyond $z=4.3$) by Ghisellini et al. (2010a). In the latter paper the observational constrain on the blazar density with black holes heavier than $10^9 M_\\odot$ in the redshift bin $5<z<6$ was based on the detection of only one object: Q0906+6930. Assuming it was the only blazar in the entire sky in this redshift bin, Ghisellini et al. (2010a) derived a comoving density of $2.63\\times 10^{-3}$ Gpc$^{-3}$ of blazars hosting an heavy black hole in the $5<z<6$ bin (see Fig. 15 in that paper).  SDSS 1146+403 was selected in the SDSS catalog covered by FIRST observations, and the common area of the sky of these two surveys is 8770 square degrees. It is the second source that can be classified as a blazars in this SDSS+FIRST survey (together with B2 1023+25). The comoving volume in the redshift bin $5<z<6$ is 380 Gpc$^3$. Therefore the number density of blazars in this redshift bin is $N_{BL}= 2\\times (40,000/8,770)/ 380 = 2.4\\times 10^{-2}$ Gpc$^{-3}$. Both B2 1023+25 and SDSS 1146+403 have black holes with $M>10^9 M_\\odot$. To find out the density of black holes in jetted sources heavier than one billion solar masses, we should consider that for each jet observed within a viewing angle $\\theta_{\\rm v}=1/\\Gamma$, there exist another $2\\Gamma^2$ sources pointing in other directions. So the number density of black holes with $5<z<6$, with a mass exceeding $10^9M_\\odot$, is $0.024 \\times 2 \\times 169 (\\Gamma/13)^2=8.1$ Gpc$^{-3}$. Multiplying by the comoving volume gives $\\sim$3,000 heavy black holes in jetted sources only.  As mentioned in the introduction, the finding of heavy and early black holes in sources with jets can severely challenge our understanding of black hole growth, especially if we associate the presence of the jet with a rapidly spinning black hole. A Kerr black hole is in fact more efficient to transform gravitational energy into radiation than a non spinning (radio--quiet) black hole. This lends support to the possibility of super--Eddington accretion (Volonteri \\& Silk 2014), and/or to the possibility that part of the gravitational energy of the accreting matter is not used to heat the disk, but to amplify the magnetic fields necessary to extract the rotational energy of the black hole (see e.g. Jolley \\& Kunzic 2008; Shankar et al. 2008, Ghisellini et al. 2013 )."
    },
    "1402/1402.1165.txt": {
        "abstract": "{A large-scale hydrodynamical cosmological simulation, \\hagn, is used to investigate the alignment between the spin of galaxies and the cosmic filaments above redshift 1.2. The analysis of more than 150 000 galaxies per time step in the redshift range $1.2<z<1.8$ with morphological diversity shows that the spin of low-mass blue galaxies is preferentially aligned with their neighbouring filaments, while high-mass red galaxies tend to have a perpendicular spin. The reorientation of the spin of massive galaxies is provided by galaxy mergers, which are significant in their mass build-up. We find that the stellar mass transition from alignment to misalignment happens around $3\\times 10^{10}\\, \\rm M_\\odot$. Galaxies form in the vorticity-rich neighbourhood of filaments, and migrate towards the nodes of the cosmic web as they convert their orbital angular momentum into spin.  The signature of this process can be traced to the properties of galaxies, as measured relative to the cosmic web. We argue that a strong source of feedback such as active galactic nuclei is mandatory to quench \\emph{in situ} star formation in massive galaxies and promote various morphologies. It allows mergers to play their key role by reducing post-merger gas inflows and, therefore, keeping spins misaligned with cosmic filaments.} ",
        "introduction": "Theoretical models of structure formation by gravitational instability and numerical simulations have predicted that small fluctuations from the early Universe lead to the formation of a large-scale cosmic web made of clustered haloes, filaments, sheets and voids~\\citep[e.g.][]{zeldovichetal82,klypin&shandarin83, Blumenthaletal84, Davisetal85}. The resulting properties of the Universe's large-scale structure are the interplay of the planar  local collapse, as emphasized in~\\cite{zeldovich70} \\citep[see also][]{shandarin&zeldovich89} and the inherent structure of the  Gaussian initial density and velocity shear fields,  leading to the cosmic web picture of  dense peaks connected by filaments,  framing the  honeycomb-like structure of walls~\\citep{bbks,bkp96}.  The extension of the Center for Astrophysics redshift survey~\\citep{huchraetal83} gave spectacular observational evidence~\\citep{delapparentetal86,Gelleretal89} for this picture, triggering a renewed interest for such large-scale galaxy surveys~\\citep{collessetal01, tegmarketal04}.  Modern simulations have established a tight connection between the geometry and dynamics of the large-scale structure of matter, on the one hand, and the evolution of the physical properties of forming galaxies, on the other. Observational information on the morphology of galaxies and its dependence on environment is routinely becoming available for galaxies up to redshift two and beyond \\citep{abrahametal07, oeschetal10, leeetal13}. Matched samples at low and high redshifts allow for the study of the evolution of many physical properties of galaxies for most of the history of our Universe in unprecedented detail. A key question formulated decades ago is nevertheless not satisfactorily answered: what properties of galaxies are driven by the cosmic environment?  There is ample literature~\\citep[e.g.][]{hoyle49, peebles69, doroshkevich70, white84, schaefer09} on the tidal torque theory (TTT). It aims to explain the early acquisition of the spin\\footnote{Hereafter, the spin is the angular momentum unit vector for simplicity.} of haloes, in the regime where the dynamics is well described by the Zel'dovich approximation, and when it is legitimate to assume that the tidal and the inertia tensors are uncorrelated. Within this framework, TTT predicts that the spin of haloes should be perpendicular to the direction of filaments.  Using \\emph{N}-body simulations, which model only dark matter,~\\cite{hahnetal07} and~\\cite{zhangetal09} found halo spins preferentially oriented perpendicular to the filaments independent of halo mass. \\citet{hatton&ninin01} claimed to detect an alignment between spin and filament, while \\citet{faltenbacheretal02} measured a random orientation of the spins of haloes in the plane perpendicular to the filaments. More recently a consensus  emerged when several works \\citep{aubertetal04,bailin&steinmetz05, calvoetal07, hahnetal07b, pazetal08, sousbie08, libeskind12, trowlandetal13} reported that large-scale structures - filaments and sheets - influence the direction of the angular momentum of haloes in a way originally predicted by~\\cite{sugermanetal00} and \\cite{lee&pen00}. These studies pointed towards  a mass-dependent orientation of the spin, arguing for the first time that the spin of high-mass haloes tends to lie perpendicular to their host filament, whereas low-mass haloes have a spin preferentially aligned with it. Nevertheless, the detected correlation remained weak and noisy until~\\cite{codisetal12} confirmed it. They quantified a redshift-dependent mass transition $M_{\\rm tr, h}$, separating aligned from perpendicular haloes, and interpreted the origin of the transition in terms of large-scale cosmic flows. \\cite{codisetal12} found that high-mass haloes have their spins perpendicular to the filament because they are the results of mergers, a scenario suggested earlier by~\\cite{aubertetal04} \\citep[see also][]{bailin&steinmetz05}. Low-mass haloes are not the products of mergers and acquire their mass by gas accretion in the vorticity-rich neighbourhood of filaments, which explain why their spins are initially parallel to the filaments~\\citep{laigle2014, Libeskind13a}.  \\cite{Tempel13} recently found tentative evidence of such alignments in the Sloan Digital Sky Survey (SDSS) with an orthogonality for elliptical galaxies and a weak alignment for spiral galaxies~\\citep[see also][]{tempel&libeskind13}. \\cite{zhangetal13} found that the major axis of red galaxies is parallel to their host filaments and is the same for blue galaxies albeit with a weaker signature. Similar measurements have been done for galaxies and walls; there is evidence that the spin of galaxies also lies within the walls in which they are contained~\\citep{trujilloetal06}.  Besides those attempts to relate the spins of galaxies with the cosmic structure, much observational effort has been made to control the level of intrinsic alignments of galaxies as a potential source of systematic errors in weak gravitational lensing measurements~\\citep[e.g.][]{C+M00,HRH00,H+S04}. Such alignments are believed to be the major source of systematics of the future generation of lensing surveys like \\emph{Euclid} or Large Synoptic Survey Telescope (LSST). Direct measurements of the alignment of the projected light distribution of galaxies in wide-field imaging data seem to agree on a contamination at a level between a few per cent and $\\sim 10$ per cent of the shear correlation functions, although the amplitude of the effect depends on the population of galaxies considered~\\citep[e.g.][]{L+P02,Joa++13a}. Given this dependence, it is difficult to use DM-only simulations as the sole resource to predict and control intrinsic alignments despite some success with the addition of a semi-analytical model prescription \\citep[e.g.][]{Joa++13a}. The inherently anisotropic nature of the large-scale structure and its complex imprint on the shapes and spins of galaxies may prevent isotropic approaches from making accurate predictions.  Very few attempts have been made to probe the degree of correlation between \\emph{galaxy} spins and their embedding cosmic web using hydrodynamical cosmological simulations. \\cite{hahn10} simulated the vicinity of a large-scale cosmic filament and found that -- at odds with the results presented in our paper -- the spin of galaxies within high-mass haloes is aligned with the filament while the spin of galaxies in low-mass haloes is perpendicular to the filament. \\cite{gayetal10} focused on the colour gradients relative to the cosmic web using the \\mn simulation~\\citep{devriendtetal10}  which did not display much morphological diversity. They found evidence of metallicity gradients towards and along the filaments and nodes of the cosmic web. \\cite{danovichetal11} also studied the feeding of massive galaxies at high redshift through cosmic streams using the \\mn simulation. They found that galaxies are fed by one dominant stream (with a tendency to be fed by three major streams), streams tend to be co-planar (in the stream plane), and that there is a weak correlation between spin of the galaxy and spin of the stream plane at the virial radius, which suggests an angular momentum exchange at the interface between streams and galaxies (see also~\\citealp{tillson12}).  In this paper, our focus is on the influence of the cosmic web as an anisotropic vector of the gas mass and angular momentum which ultimately shape galaxies. Our purpose is to determine if the mass-dependent halo spin--filament correlations of~\\cite{codisetal12} can be recovered via the morphology and physical properties of \\emph{simulated galaxies}. We aim to test these findings on a state-of-the-art hydrodynamical simulation, the so-called \\hagn simulation, which produced over 150 000 resolved galaxies displaying morphological diversity by redshift $z=1.2$, in order to identify the effect of the environmentally driven spin acquisition on morphology and to probe the tendency of galaxies to align or misalign with the cosmic filaments as a function of galactic properties.  The paper is organized as follows. Section~\\ref{section:virtual} describes the numerical set-up of our simulation, the post-processing of galaxy properties and filament tracing. Section~\\ref{sec:result} presents the excess probability of alignment as a function of morphological tracers and investigates its redshift and spatial evolution. Section~\\ref{sec:merger} discusses the origin of the observed (mis)alignment between galactic spins and filaments. We finally conclude in Section~\\ref{sec:conclusion}.  %=============================================== %  setup %=============================================== ",
        "conclusions": "\\label{sec:conclusion}  Our analysis of the \\hagn flagship simulation, which models AGN as well as stellar feedback so as to produce morphological diversity, shows that the orientation of the spin of galaxies depends on various galaxy properties such as stellar mass, $V/\\sigma$, sSFR, $M_{20}$, metallicity, colour and age. The spins of galaxies tend to be preferentially parallel to their neighbouring filaments, for low-mass, young, centrifugally supported, metal-poor, bluer galaxies, and perpendicular for higher mass,  higher velocity dispersion, red, metal-rich old galaxies. The alignment is the strongest, the closer to the filaments and further from the nodes of the cosmic web the galaxies are. This is in agreement with the predictions of~\\cite{codisetal12} for DM haloes. We find a transition mass, $M_{\\rm tr, s} \\simeq 3\\times 10^{10}\\, \\rm M_\\odot$ which is also consistent with these authors' predictions for the corresponding halo transition mass. Due to the weak galaxy--halo alignment, the amplitude of the correlation with cosmic filaments is somewhat weaker for galaxies than for haloes. It also decreases with cosmic time due to mergers and quenching of cold flows and star formation. Hence, our results suggest that galaxy properties can be used to trace the spin swings along the cosmic web.  The transition from the aligned to the misaligned case is dynamically triggered by mergers (the frequency of which increases with galaxy mass) that swing the spin of galaxies. AGN feedback has a key role at preventing further gas inflow and quenching the \\emph{in situ} star formation after such an event, in order to preserve the misalignment operated by the merger.  Though it is expected that galaxy modelling will improve over the next decade -- in particular the way feedback is taken into account in large-scale cosmological simulations --  we do not anticipate that the particular results found  in this paper should  change qualitatively. The finding that the morphological diversity of galaxies traces  populations with different spin--filament alignments, which in turn is in part inherited from the anisotropy of the embedding cosmic web, is both robust predictions of the current gravitational clustering scenario and of our understanding of  the dynamics of elliptical and spiral galaxies.  In a forthcoming paper, we will analyse more realistic mock catalogues from the \\hagn light-cone down to a lower redshift to investigate the amount of modification of the signal induced by dust extinction, projection effects, limited resolution and finite signal-to-noise ratio. More efficient and robust estimators for morphology, either intrinsic using the full data set of the simulation or projected using virtual degraded observables, will be built and compared. Quantitative comparisons to observations are postponed to this paper.  %============================================== % Acknowledgments %=============================================="
    },
    "1402/1402.5542_arXiv.txt": {
        "abstract": "Near infrared (NIR) flare/rebrightening in the afterglow of the short hard gamma ray burst (SHB) 130603B measured with the Hubble Space Telescope (HST) and an alleged late-time X-ray excess were interpreted as possible evidence of a neutron-star merger origin of this SHB. However, the X-ray afterglow that was measured with the Swift-XRT and Newton XMM have the canonical behaviour of a synchrotron afterglow produced by a highly relativistic jet. The H-band flux observed with HST 9.41 days after burst is that expected from the measured late-time X-ray afterglow. A late-time flare/re-brightening of a NIR-Optical afterglow of SHB can be produced by jet collision with an interstellar density bump, or by a kilonova, but jet plus kilonova can be produced also by the collapse of compact stars (neutron star, strange star, or quark star) to a more compact object due to cooling, loss of angular momentum, or mass accretion. ",
        "introduction": "Stripped envelope supernova explosions and neutron star mergers in close binaries were originally suggested by Goodman, Dar and Nussinov~(1987) as possible sources of cosmological gamma ray bursts. However, their proposed underlying mechanism - a spherical fireball produced by neutrino-antineutrino annihilation into electron positron pairs beyond the surface of the collapsing/merging star- turned out not to be powerful enough to produce GRBs observable at very large cosmological distances. Consequently, Shaviv and Dar~(1995) proposed that highly relativistic jets of ordinary matter are probably ejected in such events, and produce narrowly collimated GRBs by inverse Compton scattering of circumstellar light. They also suggested that short GRBs may also be produced by highly relativistic jets ejected in the phase transition of compact stars, such as neutron stars, strange stars and quark stars, into more compact objects due to mass accretion or to cooling and loss of angular momentum via winds and radiation. After the discovery of GRB afterglows, Dar (1998) proposed that they are highly beamed synchrotron radiation produced in the collision of these highly relativistic jets with the interstellar matter.  By now, there is convincing evidence that long duration GRBs are produced mostly by highly relativistic jets launched in stripped envelope supernova explosions (mainly of type Ic), but, despite enormous observational efforts, the origin of short duration GRBs remains unknown. In fact, the circumstantial evidence that has been claimed to link short hard GRBs (SHBs) with neutron star mergers in close binaries, such as their location in both spiral and elliptical galaxies and the distribution of their location offsets relative to the center of their host galaxies, which extend to a distance of ~100 kpc (e.g., Berger et al.~2013 and references therein) and beyond? (e.g., SHB 080503 with the lack of a coincident host galaxy down to 28.5 mag in deep Hubble Space Telescope imaging, Perley et al.~2009)  actually favours single compact stars with large natal kick velocity as progenitors of SHBs (Dado, Dar and De R\\'ujula 2009a) over neutron star binaries, which have much lower Galactic velocities.    It was predicted that neutron star mergers create significant quantities of neutron-rich radioactive nuclei whose decay should result in a faint transient in the days following the burst, a so-called kilonova or macronova (Li and Paczynski 1998). Recently, the broad band afterglow of the SHB 130603B (Melandri et al.~2013; Golenetskii et al.~2013) that was measured with the Swift X-ray telescope (XRT), Newton XMM, HST and ground based optical and radio telescopes was interpreted as possible evidence of a neutron-star merger origin of SHB 130603B (Tanvir et al.~2013a,b; Berger et al.~2013; Fong et al.~2013). However, in this letter we show that the X-ray afterglow of SHB 130603B, which was measured with Swift XRT (Swift/XRT GRB light-curve repository, Evans et al.~2009) and Newton XMM (Fong et al.~2013) has the canonical behaviour of a synchrotron afterglow produced by a highly relativistic jet propagating in a normal interstellar environment, as predicted by the cannonball model of GRBs (Dado, Dar and De R\\'ujula 2002, 2009a,b; Dado and Dar~2013) long before its empirical discovery (Nousek et al.~2006). This canonical X-ray afterglow does not have a \"mysterious late-time X-ray excess\" as claimed in Fong et al.~2013, and the flux observed in the NIR H-band with HST 9.6 days after burst (Tanvir et al.~2013a,b) is that expected from the measured late-time X-ray afterglow.  Moreover, a fast decline of a late-time afterglow followed by a re-brightening/flare in the NIR and optical afterglow of a GRB can be produced by a jet colliding with a density bump in the interstellar medium (e.g., Dado, Dar and De R\\'ujula 2002, 2009b), as was observed in several long-duration GRBs, such as 030329 (Lipkin et al. 2004; Matheson et al.~2003) 070311  (Guidorzi et al.~2007) and SHBs such as 050724 (Malesani et al.~2007) and 080503 (Perley et al.~2009). The host galaxy of SHB 130603B at redshift ${\\rm z=0.356}$ (Thone et al.~2013), as seen in high-resolution HST imaging (Tanvir et al.~2013b), is a perturbed spiral galaxy due to interaction with another galaxy. SHB 130603B was located  in a tidally disrupted arm. The interaction of the SHB jet with such a bumpy environment may have caused the flare/re-brightening in the NIR afterglow observed with the HST on day 9.41.  Furthermore, a late-time flare/re-brightening of a NIR-Optical afterglow of SHB can be produced by either a jet collision with an interstellar density bump or a kilonova, but jet plus kilonova can be produced also by collapse of compact stars (neutron star, strange star, or quark star) to a more compact object due to cooling, loss of angular momentum, or mass accretion. In fact, single compact stars are more likely to be found at large offsets from galactic center/disks (e.g., in the galactic halos or the intergalactic medium) than neutron star binaries or black hole-neutron star binaries. ",
        "conclusions": "Several explanations of the re-brightening of the NIR afterglow of SHB 130603B around 9.41 days after burst have been proposed. These include a macronova/kilonova produced by a neutron star merger in a close binary due to gravitational wave emission (Tanvir et al.~2013b; Fong et al.~2013), an active millisecond magnetar produced in a neutron star merger in close binaries (Metzger and Piro 2013), and a late-time flare produced by collision of the SHB jet with an ISM density bump (Dado, Dar and De R\\'ujula 2002, 2009a).  A jet plus a mini-supernova/macronova/kilonova, however, are not unique to the neutron star merger scenario. They can be produced also in a phase-transition/collapse of compact stars (neutron star, strange star or quark star) to a more compact object due to cooling, loss of angular momentum or mass accretion.  The X-ray afterglow that was measured with the Swift-XRT and Newton XMM has the expected canonical behaviour of a synchrotron afterglow produced by a highly relativistic jet. Its late time behaviour does not support a milisecond magnetar as the power-source of the chromatic afterglow of SHB 130603B. The late-time H-band flux observed with HST 9.41 days after burst is that expected from an ordinary synchrotron radiation from a jet that produced the measured late-time X-ray afterglow.  Late-time flare/re-brightening of a NIR-Optical afterglow of an SHB can be produced also by jet collision with an interstellar density bump, as seen in several GRBs. The host galaxy of SHB 130603B as seen in high-resolution HST imaging (Tanvir et al.~2013b) is a perturbed spiral galaxy due to interaction with another galaxy. The GRB was located in in a tidally disrupted arm of its host galaxy. The interaction of the GRB jet with such a bumpy environment could produce the flare/re-brightening of the NIR afterglow of the GRB observed with the HST 9.41 days after burst.  The star formation within the host, location of SHB 130603B on top of the tidally disrupted arm, strong absorption features and large line of sight extinction that were observed indicate that the GRB progenitor was probably not far from its birth place (de Ugarte Postigo et al.~2013), untypical to the usually long life time before neutron star merger due to gravitational wave emission in the known neutron star binaries in our galaxy. Moreover, the failure to detect a host galaxy down to 28.5 mag in deep Hubble Space Telescope imaging searches in the case of e.g., SHB 080503 with a late time flare/rebrightening (Perley et al.~2009) suggest a large natal kick velocity of its progenitor, unlikely for compact binaries, but often observed for isolated neutron stars/pulsars  (Hobbs et al.~2005).  The true smoking gun for the neutron star merger in close binaries is the detection of gravitational waves. Unfortunately, this is unlikely to occur before the completion of the new generation of gravity-wave detectors, as the sensitivity of current detectors such as LIGO, and Virgo, is several orders of magnitude below what would be required to detect a merger at a distance similar to the nearest SHBs with known redshift.   {\\bf Acknowledgement: We thank  A. de Ugarte Postigo for useful comments.}"
    },
    "1402/1402.1227_arXiv.txt": {
        "abstract": "{In this paper we will illustrate how to constrain unavoidable K\\\"ahler corrections for ${\\cal N}=1$ supergravity (SUGRA) inflation from the recent Planck data. We will show that the non-renormalizable K\\\"ahler  operators will induce in general {\\it non-minimal kinetic} term for the inflaton field, and two types of SUGRA corrections in the potential - the {\\it Hubble-induced mass} ($c_{H}$), and the {\\it Hubble-induced A-term} ($a_{H}$) correction. The entire SUGRA inflationary framework can now be constrained from (i) the {\\it speed of sound}, $c_s$, and (ii) from the upper bound on the {\\it tensor to scalar ratio}, $r_{\\star}$. We will illustrate this by considering  a heavy scalar degree of freedom at a scale, $M_s$, and a  light inflationary field which is responsible for a slow-roll inflation. We will compute the corrections to the kinetic term and the potential for the light field explicitly. As an example, we will consider a visible sector inflationary model of inflation where inflation occurs at the point of {\\it inflection}, which can match the density perturbations for the cosmic microwave background radiation, and also explain why the universe is filled with the Standard Model degrees of freedom. We will scan the parameter space  of the non-renormalizable K\\\"ahler operators, which we find them to be order ${\\cal O}(1)$, consistent with physical arguments. While the scale of heavy physics is found to be bounded by the tensor-to scalar ratio, and the speed of sound,  $ {\\cal O}(10^{11}\\leq M_s\\leq 10^{16})$~GeV, for $0.02\\leq c_s\\leq 1$ and $10^{-22}\\leq r_\\star \\leq 0.12$. }    \\begin{document} ",
        "introduction": "The success of primordial inflation~\\cite{Guth:1980zm}, for a review, see~\\cite{Mazumdar:2010sa}, can be gauged by the current observations arising from the comic microwave background (CMB) radiation~\\cite{Planck-infl,Planck-1,Ade:2013ydc}. The observations from Planck have put interestingly tight bounds on a number of unknown parameters of a generic inflationary model~\\cite{Planck-infl}, in particular the {\\it speed of sound}, $c_s$, of the perturbations, which also determines any departure from the Gaussian perturbations, the local type of non-Gaussianity, $f_{NL}^{\\mathrm{local}}$,  and the constraint on {\\it tensor-to-scalar ratio}, $r_\\star$, which can potentially unearth the scale of New Physics within any given {\\it effective field theory } set-up.  The ${\\cal N}=1$ supergravity (SUGRA)~\\cite{Freedman:1976xh}, for a review, see~\\cite{Nilles:1983ge}, is an excellent well-defined set-up where we can address some of the key questions about the physics of the new scale for instance. The K\\\"ahler metric determines the kinetic term for the inflaton  potential, and one particular choice is the {\\it minimal} Kinetic term for the inflaton field. However, quantum corrections to the K\\\"ahler potential is not very well-known. Generically corrections to the K\\\"ahler potential arise from integrating out the heavy physics, and due to lack of concrete knowledge on the details of the heavy physics, many times computing these corrections can be very challenging, see~\\cite{Berg:2005ja}.  The aim of this paper is to place a generic bound on the Planck suppressed corrections to the K\\\"ahler potential on top of the minimal K\\\"ahler potential. We will consider dimensional $3$ and dimensional $4$ {\\it gauge invariant} non-renormalizable K\\\"ahler operators in this paper. Since these correction will lead to a departure from the minimal kinetic term for the inflationary potential, such corrections can now be bounded from the Planck data, especially from the speed of sound of the primordial perturbations. In fact the K\\\"ahler potential within ${\\cal N}=1$ SUGRA can induce corrections to the inflationary potential which can yield large Hubble-induced mass correction to the inflaton field~\\cite{Bertolami:1987xb,Mazumdar:2011ih}, some times known as the SUGRA-$\\eta$ problem, and the Hubble-induced SUGRA $A$-term for the potential.  In the context of ${\\cal N}=1$ SUGRA hybrid inflation~\\cite{Rehman:2010wm,Khalil:2010cp,Shafi:2010jr}, some of the K\\\"ahler corrections were constrained by the tensor to scalar ratio, $r_\\star$, and the spectral tilt, $n_S$, of the power spectrum. In this paper we will be constraining these K\\\"ahler corrections systematically from the interaction of the heavy physics to the light inflaton field from the recent Planck data. Let us consider the scale of New Physics, $M_s$, to be within \\begin{equation} M_p\\geq M_s\\geq H_{inf}\\,, \\end{equation} where $H_{inf}$ is the Hubble parameter during inflation, and $M_p=2.4\\times 10^{18}$~GeV. In order to constrain $M_s$, we would require to consider at least $2$ fields, one which is heavy at the relevant scale, $M_s$,  and the other which is light. We will assume that these two fields are coupled gravitationally. In past such a scenario has been considered by many authors, where the heavy field leaves interesting imprints in the dynamics of a low scale inflation~\\cite{Burgess:2002ub,Burgess:2003zw,Burgess:2012dz,Achucarro:2012yr,Achucarro:2012sm,Collins:2009pf,Achucarro:2010da,Achucarro:2010jv,Assassi:2013gxa,Agarwal:2013rva,Antusch:2008pn}. Broadly speaking there are two possible scenarios which one can envisage:  \\begin{itemize}  \\item {\\bf The heavy field is dynamically frozen}: we can imagine that the heavy field is completely {\\it frozen}, in which case it would be effectively a single light field with  a canonical kinetic term for the inflaton field, with a speed of sound, $c_s=1$. For a slow roll inflation, the perturbations will be primarily Gaussian. If the heavy field is settled down to its minimum VEV, i.e. zero, then there will be no effect from the heavy field at all. However, if the heavy field is settled with a finite non-zero VEV, and it remains dynamically inactive, means its velocity is strictly zero, then it can still contribute to the vacuum energy density of the inflaton, and this would be encoded in $H_{inf}$. Also, the kinetic term for the light inflaton field will depart from being pure canonical. However the departure will depend on the scale of new physics. If $M_s\\ll M_p$, then the departure from canonical kinetic term will be negligible for all practical purposes. Therefore, again the observational predictions for the CMB will be unaltered and will be similar to the previous case. Both of these scenarios were taken into account by various interesting papers, see for example~\\cite{Rehman:2010wm,Shafi:2010jr,Mazumdar:2011ih,Antusch:2010mv,Antusch:2012jc,Choudhury:2013jya}, and here we will not consider them in great details. We will analyse a slightly different scenario as mentioned below.     \\item{\\bf The heavy field is coherently oscillating during the initial phases of  inflation}:  in this case we will consider a very simple scenario, where we imagine that the heavy field is coherently oscillating at a VEV given by $M_s$ with an amplitude $M_s$ at the onset of inflation driven by the light field. The coherent oscillations of the heavy field will not last forever, its amplitude would be damped during inflation very rapidly within couple of e-foldings of inflation. However, just right at the onset of inflation, the relevant modes which are leaving the Hubble patch for the CMB can be constrainable. This will provide a window of opportunity for us to constrain such a scenario, see Refs.~\\cite{Burgess:2002ub,Burgess:2003zw,Burgess:2012dz} for probing the influence of heavy physics into the light inflaton field. In this paper we will consider a similar scenario, but in the context of ${\\cal N}=1$ SUGRA.  First of all the coherent oscillations of the heavy field around its non-zero vacuum would provide a non-zero vacuum energy density, i.e $\\sim M_s^4$. Through its coupling to the light field in the K\\\"ahler potential, it would also yield non-canonical kinetic term contribution to the light field, and therefore $c_s\\neq 1$ for a slow roll inflaton field. Eventually, the heavy field will be settled down to its minimum. We presume that the dominant contribution to the long wavelength fluctuations are still seeded by the light  inflaton, but the fact that  $c_s\\neq1$ for the inflaton, it would leave imprints which would be constrainable directly by the scale of heavy physics, $M_s$, and the non-minimal corrections to the k\\\"ahler potential. \\end{itemize}  We will discuss this latter scenario in some details, and provide a full ${\\cal N}=1$ SUGRA potential for the light and the heavy field within a simple example. We will be using the following constraints from CMB, and also requirement for a guaranteed reheating of the Standard Model d.o.f for the success of big bang nucleosynthesis~\\cite{Fields:2006ga}.  \\begin{enumerate}  \\item{Successful single field inflation driven by $\\phi$ field with the right amplitude and tilt of the power spectrum. \\begin{eqnarray}\\label{powq1} \\displaystyle 2.092<10^{9}P_{S}<2.297~~(~within~2\\sigma),\\\\ \\label{powq2} \\displaystyle 0.958<n_{S}<0.963~~(~within~2\\sigma)\\,. \\end{eqnarray} }  \\item{Speed of sound, $c_s$: The Planck analysis has constrained it to be~\\cite{Planck-infl,Ade:2013ydc}: \\begin{equation}\\label{c-s} \\displaystyle 0.02 \\leq c_s \\leq 1~~(~within~2\\sigma)\\,. \\end{equation} } \\item{Tensor-to-scalar ratio: The Planck constraint is $r_\\star\\leq 0.12$~\\cite{Planck-infl} \\begin{equation}\\label{t-s} \\displaystyle r_\\star \\equiv \\frac{P_T}{P_S}\\leq 0.12. \\end{equation} }  \\item{Local type of non-Gaussianity, $f_{NL}^{\\mathrm{local}}$: The Planck constraint on local non-Gaussianity is~\\cite{Planck-infl,Ade:2013ydc}: \\begin{eqnarray}\\label{f-nl} \\displaystyle f_{NL}^{\\mathrm{local}}=2.7\\pm 5.8~~(~within~1\\sigma)\\,. \\end{eqnarray} In this paper we will not consider the constraints arising from various non-Gausisnaity bounds~\\cite{Ade:2013ydc}, but we will solely focus  on the constraints arising from the speed of sound during perturbation, and the tensor to scalar ratio. In the companion paper we have discussed how non-Gaussianity can constrain the K\\\"ahler corrections.}  \\item{Particle physics constraint: We wish to ensure that the inflaton {\\it solely} decays into the Standard Model (SM) d.o.f, therefore we embed the light fields within supersymmetric SM, such as minimal supersymmetric Standard Model (MSSM). In this case the inflaton carries {\\it solely} the SM charges as in the case of the inflaton driven by {\\it gauge invariant} combinations of squarks and sleptons~\\cite{Allahverdi:2006we,Allahverdi:2006iq,Allahverdi:2006cx,Chatterjee:2011qr}. This will naturally ensure that we obtain the right abundances for the dark matter and the baryons in the universe as required by the observations~\\cite{Allahverdi:2007vy,Balazs:2013cia}. One can however follow a hidden sector or a SM gauge singlet inflaton, but it is not always straightforward to explain the universe with the right dark matter abundance~\\cite{Dev:2013yza}, and baryon asymmetry, see~\\cite{Mazumdar:2010sa}.}  \\end{enumerate}  The results of the first half of this paper will be very generic - applicable to any inflationary scenario. In section \\ref{perturb}, we will discuss briefly the Planck constraints. In section \\ref{setup}, we will discuss the setup with one heavy and one light superfield which are coupled via gravitational interactions through K\\\"ahler potential. In section \\ref{eft}, we will describe the effective field theory potential for the light superfield $\\Phi$, and discuss the kinetic terms for various interesting scenarios. In section \\ref{N-M}, we will discuss the role of non-canonical kinetic term and consider two possibilities, one where the heavy superfield is dynamically frozen, see section \\ref{frozen}, and the more interesting scenario when the heavy field is oscillating at the onset of inflation, see section~\\ref{oscillating}. We will  scan the parameters for the Planck suppressed K\\\"ahler operators in subsection~\\ref{02},  we will discuss how tensor-to-scalar ratio, $r_\\star$, can constrain the mass scale of the heavy physics. ",
        "conclusions": "In this paper, we have shown that in any ${\\cal N}=1$ SUGRA inflation  model when ever there are more degrees of freedom, non-minimal K\\\"ahler corrections would induce three distinct types of corrections: (i) non-minimal kinetic term for the inflaton, (ii) Hubble-induced mass correction to the inflaton, and (iii) Hubble-induced $A$-term in the potential.  The exact nature of K\\\"ahler potential and K\\\"ahler corrections might not be known in all possible scenarios, but our aim has been to constrain the coefficients of the non-renormalizable K\\\"ahler higher dimensional operators phenomenologically, which are gauge invariant, from the recent Planck data. We assumed minimal K\\\"ahler potentials for all the fields to begin with. We first considered the heavy physics to be completely decoupled from the dynamics of the light inflaton field. We considered the light field to be embedded within MSSM, such that the reheating of the universe is guaranteed to be that of the SM dof. In the simplest setup when the heavy field is well settled down in its potential, it only affects via its vacuum energy density. The kinetic terms are mostly canonical, and therefore we do not obtain any constraint on the coefficients of the dimensional $3$ and $4$ non-reormalizable K\\\"ahler operators.  We further investigated an intriguing possibility, when the heavy field is coherently oscillating with a frequency larger than the Hubble parameter during the onset of inflation, while the light field is slowly rolling over the potential.  In this particular scenario, we were able to constrain the coefficients of the Planck suppressed K\\\"ahler operators of dimensional $3$ and $4$. We scanned the four parameters, $a,~b,~,c~,d$, and obtained a region of the parameter space where we can satisfy the current Planck observations, i.e. $P_S,~n_S,~c_s$ and $r_\\star$ within $2\\sigma$~CL, and we obtained all the coefficients to be of order $a,~b,~c,~d \\sim {\\cal O}(1)$, as naturally expected in any non-rrenormalizable SUGRA theory. In fact, as we can see from Fig.~(\\ref{fz}) their magnitudes are always less than one.  In Fig.~(\\ref{fgh}), we have shown for the range of non-renormalizable corrections, the parameter space for the allowed range of $P_S$ versus $n_S$ for the allowed range of $0.02\\leq c_s\\leq 1$.  In Fig.~(\\ref{fig2con}),  we have plotted $r_\\star$ vs. $n_S$, for $c_s=1$ and $c_s=0.02$ for the number of e-foldings, $N=50,~70$. For the range of parameter space scanned, we were able to set an upper limit on the scale of new physics from the constraints arising from $r_\\star$, which we obtained to be within $10^{11}\\leq M_s\\leq 10^{16}$~GeV. For the lower bound on $M_s$, we found $r_\\star\\sim {\\cal O}(10^{-22})$ and extremely negligible, and for the upper bound we saturated $r_\\star =0.12$. Note that the current Planck data mildly prefers lower value of the speed of sound, i.e. $c_s < 1$, this is visible from our scans and the plot on $r_\\star$ versus $n_S$, see Fig.~(\\ref{fig2con}).  Finally, we would like to mention that all the above bounds have been obtained for a very particular kind of inflation model, which is fully embedded within MSSM, the inflaton is an MSSM flat direction and inflation happens at the point of inflection with a fine tuned parameter at the inflection point is roughly one part in $10^{4}$. We chose MSSM inflation for its advantage that the dynamics can be well understood during inflation and after inflation. In particularly, we can ascertain that the universe after inflation would be filled with the SM degrees of freedom, and also the model is capable of explaining the Higgs mass constraint and the dark matter abundance, along with the constraints on the inflaton mass arising from the LHC~\\cite{Allahverdi:2007vy}. Not every model of inflation enjoys such advantages, and therefore studying this model in some details along with SUGRA corrections yielded interesting constraints. Our methodology can be followed for other kinds of inflationary models too.  There is a further scope of improvement in our analysis. So far we have only used the Planck constraints from the power spectra, $P_S$, spectral tilt, $n_S$, tensor-to-scalar ratio, $r_\\star$, and the constraint on the speed of sound, $c_s$. In principle we should be able to use the non-Gaussian parameters, $f_{NL}^{local}$, $g_{NL}^{local}$ and possibly $\\tau_{NL}^{local}$, to further constrain the non-renormalizable K\\\"ahler operators of dimension $3$ and $4$. In our companion paper, we would consider the non-Gaussian constraints in some details. All these cosmological constraints arising from Planck and future CMB missions can further improve our understanding of many different aspects of physics beyond the SM. With an improvement on tensor-to-scalar ratio, $r_\\star$, we would be able to further constraint the scale of heavy physics, $M_s$."
    },
    "1402/1402.4772_arXiv.txt": {
        "abstract": "The Temporal and Spectral variations of black hole candidates during outbursts have been reported in several publications. It is well known that during an outburst, the source becomes soft in the first few days, and then returns to the hard state after a few weeks or months. In the present paper, we show the variation of Comptonization Efficiency (CE), obtained from the ratio of the black body photon number to the power-law photon number, as a function of time in several outbursts. Since the power-law photons are generated through inverse-Comptonization of the intercepted soft photons, the CE is a measure of the geometry of the Compton cloud. Our investigation indicated that all the outbursts starts with a large CE and becomes very small after a few days, when the Compton cloud becomes very small to intercept any significant number of soft photons. CE returns back to a larger value at the end of the outburst. ",
        "introduction": "\\label{s:intro}  It is well known that the  black hole accretion flows in a compact binary system typically consist of a Keplerian disk which emits soft or low energy photons and a hot Compton cloud which inverse Comptonizes the soft photons into high energy photons. When the black hole spectral state changes from the hard state to the soft state and vice versa \\citep{C95} the Compton cloud must change its shape and temperature: In the soft state when the accretion rate in the Keplerian disk is high, the Compton cloud is smaller and cooler, while in the hard state, when mass accretion rate via Keplerian disk is very low, the Compton cloud is larger and hotter. There are various models of the Compton cloud in the literature, ranging from the hot Corona, to post-shock region of an accreting low angular momentum flow. Outburst sources are ideal candidates to study the changes in the size of the Compton cloud. This is because, the object is known to be in the hard state at the beginning of the outburst, but in a matter of few days to a few weeks, it changes its states to other states, thereby giving us a unique opportunity to study the size of the Compton cloud very well. Fig.~\\ref{outgeo} shows the cartoon diagram of variation of CENBOL geometry during outburst of XTE J1550-564 \\citep{C09}. From the spectral analysis, it appears that the following sequence is typically followed by most, if not all, the outburst sources: hard $\\rightarrow$ hard-intermediate $\\rightarrow$ soft-intermediate $\\rightarrow$ soft $\\rightarrow$ soft-intermediate $\\rightarrow$ hard-intermediate $\\rightarrow$ hard) \\citep{C08, C09}. Recently, \\citet{P11,P12a} showed that the Compton cloud in the variable source GRS 1915+105 changes its size as it transits from one variability class to another. We computed a quantity called the Comptonizing Efficiency (CE) which is the ratio of the power-law photons to the blackbody photons in the spectrum at a given instant. We show that CE is very small in softer classes and larger for the harder classes. Within some of the classes, there are evidences of rise and fall of the count rates and the spectral slopes in a matter of a few seconds. Since the number of hard photons depends on the the optical depth of the Compton cloud, we clearly see the change in the optical depth of the Compton cloud in that time scale.  In the present paper, we analyze RXTE data to study the variation of the size of the Compton cloud during the outbursts of several black hole candidates. We clearly show that the outburst starts with a large Comptonization efficiency, i.e., a large sized Compton cloud with a poor soft photon source. In the rising phase, the cloud becomes smaller and smaller on a daily basis till it became minimum when the object went to a soft state. In the declining phase of the outburst, the trend is reversed. The black hole candidates we analyzed are GRO J1655-40 and XTE J1550-564.  \\begin{figure} \\begin{center} \\begin{tabular}{p{5.5cm}cp{5cm}} \\raisebox{-\\height} {\\includegraphics[viewport=0 150 1000 810,clip=,width=7.0cm]{partha_01.eps}} & \\quad & \\caption{Cartoon diagram of CENBOL geometry variation during outburst of XTE J1550-564. \\label{outgeo}} \\end{tabular} \\vskip -1.0cm \\end{center} \\vskip -1.6cm \\end{figure} ",
        "conclusions": "In this paper, we can come to a conclusion that generally speaking, all the outbursts start and end with a high Compton cloud size, though not necessarily the highest optical depth. As the outburst progresses, the CE becomes smaller at the peak of the outburst CE has the minimum value, i.e., the size of the CENBOL is very small. Since CE is a concept where the hardness ratio is determined from ratio of the dynamically obtained black body and power-law photons, it is insensitive to the mass of the black hole. This is unlike the usual hardness ratio, where photons in each energy band could be a mixture of the black body and power-law photons. We find that CE in outbursts sources may vary from almost $\\sim 0.0$ to $\\sim 3$. In contrast, the variable source GRS 1915+105 shows CE between $0.005$ to $0.8$ and has relatively high luminosity (even after factoring out the effect of the mass of the black hole), suggesting that it is in the soft-intermediate and hard-intermediate states just after the peak of a possible outburst \\citep{P12a}. If so, in future, this source may slow down its activities when the viscosity in the system is reduced. From \\citet{P12b}, we see that the range of CE is different for different objects. The duration of the outbursts, the QPO frequency range etc. are also found to be different. There is also a considerable scatter in CE which may be triggered by outflows from the CENBOL. All these require a more thorough analysis for a unifying understanding of the outbursts. This will be addressed in our future publications. \\begin{center} \\vskip -0.25cm {\\bf \\large Acknowledgments} \\end{center} \\vskip -0.25cm P. S. Pal acknowledges the SNBNCBS-PDRA fellowship. \\vskip -0.25cm"
    },
    "1402/1402.4291_arXiv.txt": {
        "abstract": "In this paper it is shown that in non-minimally coupled Brans-Dicke theory containing a self-interacting potential, a suitable conformal transformation can automatically give rise to an interaction between the normal matter and the Brans-Dicke scalar field. Considering the scalar field in the Einstein frame as the quintessence matter, it has been shown that such a non-minimal coupling between the  matter and the scalar field can give rise to a late time accelerated  expansion for the universe preceded by a decelerated expansion for very high values of the Brans-Dicke parameter $\\omega$. We have also studied the observational constraints on the model parameters considering the Hubble and Supernova data. ",
        "introduction": "Recent observational data from the Supernova Cosmology Project \\citep{1, 11, 12, 13} and the WMAP data \\citep{2, 21, 22, 23, 24} have strongly confirmed that we live in an accelerating universe. According to these observations, nearly 70$\\%$ of the total energy density of the universe has a large negative pressure which is apparently unclustered and is dubbed as Dark energy (DE), which will in fact cause the cosmic  expansion to speed up. A number of candidates have appeared in the literature which can generate the cosmic acceleration very effectively. For an excellent review on dark energy models, one can look into \\citep{3, 31, 32} and \\citep{4}. Out of all these candidates, one of the most popular candidates is a scalar field with a positive potential which can generate an effective negative pressure if the potential term dominates over the kinetic term. This scalar field is often dubbed as {\\it ``quintessence scalar field\"}. \\\\  However, quintessence scalar field models find difficulties in explaining why the energy densities of the Cold Dark Matter (CDM) and quintessence matter (Q-matter)  should be comparable today. This is known as the coincidence problem. A ``tracker field\" model can address this problem to some extent. This kind of fields has an attractor like solution with nearly independent initial conditions. However, it is not clear to what extent this field has ability to drive the current accelerated expansion of the universe as well as at the same time to give a solution to the coincidence problem.\\\\  Another way to solve the coincidence problem is to consider that the two components of matter, namely, the CDM and dark energy or the Q-matter are  interacting so that during this interaction there will be some transfer of energy from one field to another. A number of models have been proposed where the flow of energy is from the dark matter component to the dark energy component \\citep{zimdahl,zimdahl1, zimdahl2, zimdahl3, zimdahl4, zimdahl5, reddy}, so that during the late time of evolution,  the dark energy dominates over the ordinary matter and drives the acceleration. However, in most of these models, the interaction chosen is ad-hoc and such choices are not motivated by any physical theory. So, search is on for a suitable cosmologically viable interacting model of dark energy. \\\\  In this paper, we have proposed a cosmological model for describing the dynamics of such an accelerating universe in Brans-Dicke theory where the dark matter and dark energy components are shown to interact among themselves. However, the form of interaction is not chosen arbitrarily; rather it appears as a result of a particular kind of conformal transformation. These type of transformation was first proposed by Dicke \\citep{dicke} and has been extensively used since then in scalar-tensor theories of gravity \\citep{anjan,  conformal, confpapers, confpapers1, confpapers2}. The reason for choosing Brans-Dicke scalar field is that unlike minimally coupled scalar field models, this geometrical scalar field has a proper theoretical background and no one has to bother about the genesis of the scalar field. However, in original Brans-Dicke theory, the scalar field is massless; but in our model we have introduced a potential term for the scalar field. This type of scalar field have been investigated thoroughly as the driving agent for the late time acceleration of the universe \\citep{conformal, 5, 6}. The paper is organized as follows: In section 2, we have considered a non-minimally coupled Brans-Dicke (BD) scalar field, along with a self-interacting potential $V(\\phi)$. In the conformally transformed Einstein$'$s frame, we try to solve Einstein$'$s field equations and obtain different dynamical quantities in terms of the scale factor $a(t)$. In the following subsections, first we assume a specific form of the energy density and then a specific form of the scale factor and obtain the  various related parameters for both the models. In both cases, it has been shown that the deceleration parameter $q$ has a smooth transition from a positive value to some negative value which indicates that the universe was undergoing an early deceleration followed by late time acceleration which is essential for the structure formation of the universe. The important point to note here is that for both the choices, we have obtained the smooth transition from deceleration to acceleration for large values of $\\omega$, the Brans-Dicke parameter, which is consistent with the value suggested by solar system experiments \\citep{tortora}. Section 3 describes the observational constraints on various parameters of the model and the last section  discusses the results. ",
        "conclusions": "In this paper, we have shown that it is possible to build an interacting cosmological model using a non-minimally coupled scalar field. We have shown that the form of interaction chosen was not ad-hoc, rather it appeared as a result of some suitable conformal transformation. We have considered two toy models describing the evolution of the universe and in both the cases it has been found that during the evolution the universe undergoes an early phase of decelerated expansion $(q > 0)$ followed by a late time accelerated expansion phase $(q < 0)$. This is a must for the structure formation of the universe so that the universe evolves to the same form as it looks today. \\\\ In model I, we have considered a specific form for the energy density of the universe and in model II, we have considered a specific ansatz for the evolution of the scale factor of the universe. In the transformed frame, we could find out the expressions for all the relevant cosmological parameters for both the models. Although there are issues regarding the interconversion between the Einstein's frame and Jordan's frame, such as the geodesic equation is no longer valid in Einstein's frame; but things are fine as long as one sticks to any one of these frames and all the relevant parameters are defined in the same frame.\\\\ We have calculated the various parameters of the model in the transformed version of the theory. We have also shown that the results obtained are consistent with the various observational values as mentioned earlier; e.g.: in our model we have chosen the value of the parameter $W = 0.001$ which makes $\\omega > 3,00,000$ which is consistent with the solar system type experiments which puts the limit on $\\omega$ as $\\omega > 40000$ \\citep{tortora}. Also in model I, we have considered ${\\omega}_{\\phi}$ to be almost a constant ($ \\sim -0.9$ or so) which is also consistent with the values suggested by \\citep{davis, davis1}. However, for model II, we have chosen $\\alpha = 2$, $B = 0.2$ and $W = 0.001$ and we have shown that the results obtained are not vary sensitive to these values.\\\\ We have also tested our model with the observational data from the Huble and Supernova datasets. For both the models it has been found that the values of various parameters of the models which were chosen for analytical results are well fitted in the $1\\sigma$ and $2\\sigma$ confidence contours.\\\\ So, such an interacting model, where the source of interaction is not ad-hoc, can give rise to an effective solution for the {\\it{coincidence problem}} and also can provide solutions for many other cosmological problems. In this paper, we have considered two ansatz so as to obtain exact solutions for various cosmological parameters; however there are infinite number of possibilities and one can generate various effective cosmological models from such an interacting scenario which may provide even better fit to the available observational data."
    },
    "1402/1402.0547_arXiv.txt": {
        "abstract": "We present an all-sky sample of 984 candidate intermediate-mass Galactic star-forming regions color-selected from the \\emph{Infrared Astronomical Satellite} (\\emph{IRAS}) Point Source Catalog and morphologically classify each object using mid-infrared \\emph{Wide-field Infrared Survey Explorer} (\\emph{WISE}) images.  Of the 984 candidates, 616 are probable star-forming regions (62.6\\%), 128 are filamentary structures (13.0\\%), 39 are point-like objects of unknown nature (4.0\\%), and 201 are galaxies (20.4\\%).  We conduct a study of four of these regions, IRAS~00259+5625, IRAS~00420+5530, IRAS~01080+5717, and IRAS~05380+2020, at Galactic latitudes $|b|$ $>$ 5\\deg\\ using  optical spectroscopy from the Wyoming Infrared Observatory along with near-infrared photometry from the Two-Micron All Sky Survey to investigate their stellar content.  New optical spectra, color-magnitude diagrams, and color-color diagrams reveal their extinctions, spectrophotometric distances, and the presence of small stellar clusters containing 20--78 \\msun\\ of stars.  These low-mass diffuse star clusters contain $\\sim$65--250 stars for a typical initial mass function, including one or more mid-B stars as their most massive constituents. Using infrared spectral energy distributions we identify  young stellar objects near each region and assign probable masses and evolutionary stages to the protostars.  The total infrared luminosity lies in the range 190 to 960 \\lsun, consistent with the sum of the luminosities of the individually identified young stellar objects. ",
        "introduction": "} It is well established that most stars form in groups or clusters \\citep{Larson1982}.  \\citet{Weidner2006} suggested that an underlying physical relationship exists between the most massive star in an embedded cluster and the total stellar mass of the cluster.  Low-mass stars tend to form in small, loose aggregates of 10--100 stars $\\mathrm{pc}^{-3}$ \\citep{Testi1999}, while high-mass stars are associated with clusters that contain up to 1000 stars $\\mathrm{pc}^{-3}$ \\citep{Hillenbrand1995}.  Evidence suggests that there may be a turn-on mass for stellar clustering at the scale of intermediate-mass stars of 8--10 \\msun\\ and that this onset may be sudden \\citep{Testi1997} or gradual \\citep{Testi1999}.  Hence, identifying and studying clusters having intermediate-mass stars as their most massive members may yield clues about the onset of cluster formation.  While it is reasonably well understood how isolated, low-mass stars form \\citep{Mckee2007}, it may not be as simple as scaling this standard theory to explain high-mass star formation \\citep{Zinnecker2007}.  Scaling up low-mass star formation to higher masses ($>$10 \\msun) results in stars that begin to burn hydrogen while still accreting.  The resulting radiation pressure should expel the accreting material preventing more massive stars from forming, yet stars much more massive do exist \\citep{Larson1971,Zinnecker2007,Krumholz2009}.  At least two competing theories have been proposed to explain this paradox.  In the competitive accretion model, stars within a common gravitational potential accrete from a distributed gas reservoir.  The potential funnels gas down toward the cluster center, so the stars located closer to the center of the potential acrete at higher rates than isolated stars \\citep{Bonnell1997,Bonnell2001,Bonnell2006}.  In the turbulent core accretion model, molecular clumps fragment into gaseous cores that subsequently collapse to make individual stars.  As gas accretes onto the protostar, gravitational energy is released that heats up the surrounding gas.  This increase in gas temperature raises the Jeans mass which suppresses fragmentation and results in fewer, but more massive, stars.  This monolithic collapse would dictate a correlation of the clump mass with the stellar mass, as the mass of the core determines the mass reservoir available to form the star \\citep{Mckee2003,Krumholz2005,Krumholz2007}.  These divergent theories have led to increased interest in studying star-forming regions that occupy the transition from the low- to high-mass regimes.  \\citet{Arvidsson2010} defined and introduced a small sample of 70 intermediate-mass star-forming regions (\\imsfrs) --- regions producing stars up to but not exceeding $\\sim$8 \\msun.  IM SFRs populate the continuum linking the low- and high-mass star-forming regimes  where the onset of cluster formation may occur and the mechanisms that differentiate high-mass from low-mass star formation may become significant.  \\imsfrs\\ are sites of star formation distinct from regions of high-mass star formation.  They have typical luminosities of $\\sim$$10^4$ \\lsun\\ and typical diameters of $\\sim$1 pc, as judged from the sizes of the surrounding photodissociation regions (PDRs).  They are typically associated with molecular clumps of mass $\\sim$$10^{3}$ \\msun.  The molecular material associated with IM SFRs typically has an order of magnitude lower peak mass column density and clump mass compared to ultra-compact HII regions which are examples of high-mass star formation \\citep{Arvidsson2010}.  We have begun a program to study the stellar properties and the local environments surrounding \\imsfrs\\ and identify a larger more complete sample of \\imsfrs.  Our primary goal in this paper is to identify the stellar content and confirm the intermediate-mass nature of these candidate \\imsfrs.  In Section 2 we define an all-sky sample of \\imsfrs\\ and categorize them based on their mid-IR morphological appearance.  In Section 3 we use the Two Micron All-Sky Survey (2MASS) Point Source Catalog \\citep{Cutri2003}, new optical spectroscopy, and YSO SED fitting to investigate the stellar content of a sample of four \\imsfrs.  In Section 4 we present the results of our stellar content analysis, and we summarize key results in Section 5. ",
        "conclusions": "}  In this paper we morphologically classify an all-sky sample of 984 candidate \\imsfrs\\ using \\emph{WISE} mid-IR images.  The objects in this all-sky sample were selected from the \\emph{IRAS} Point Source Catalog using colors that indicate large PAH contribution and cool dust found in \\imsfrs.  We determined that 616 (62.6\\%) are classified as ``blobs/shells,'' 128 (13.0\\%) are classified as ``filamentary,'' 39 (4.0\\%) are classified as ``star-like'' objects, and 201 (20.4\\%) are classified as ``galaxies.''  We find that the blobs/shells, filamentary, and star-like objects are concentrated in the Galactic Plane, while galaxies appear to be uniformly distributed throughout the sky.  We conducted a pilot study  of four Galactic \\imsfrs\\ that are classified as blobs/shells using optical spectroscopy and near-IR photometry.  Optical spectroscopy of four--five stars per region demonstrates that each contains, within its projected boundaries, two or more intermediate-mass stars.  Eight of the eighteen spectroscopically classified stars are B-type and two are A-type with no stars more massive than $\\sim$B2 (8 \\msun) found.  Near-IR color-magnitude diagrams using 2MASS data indicate that these stars are the most massive  \\& luminous constituents, suggesting that there may be an upper limit to the mass of stars in \\imsfrs\\ of $\\sim$B2.  SEDs of candidate YSOs in each region also indicate that they are continuing to produce intermediate- and lower-mass stars, but not massive stars.  These findings support the idea that the \\emph{IRAS} color-selection of (\\citealt{Kerton2002,Arvidsson2010}) does indeed select regions producing intermediate-mass stars and that these regions can be distinguished from high-mass SFRs using \\iras\\ mid-IR colors.  Cleaned 2MASS color-magnitude diagrams reveal a strong excess of stars along the pre-main-sequence isochrones in IRAS~00420+5530 and IRAS~05380+2020 and a less obvious excess in IRAS~00259+5625 and IRAS~01082+5717, indicating the presence of small clusters associated with the spectroscopically identified IM stars.  The former two regions (IRAS~00420+5530 and IRAS~05380+2020) that have the strongest evidence for stellar clusters also the earliest stars as their most massive members ($\\sim$B2), while the latter two clusters (IRAS~00259+5625 and IRAS01082+5717) host somewhat later spectral types (B7 and B4, respectively).   Thus, the most massive stars appear to be associated with richer and more massive stellar clusters.   While the current small sample size precludes a definitive statement, this correlation is consistent with the idea that there may be a turn-on point for cluster formation at the cloud masses and stellar masses associated with the formation of mid-B spectral types as the most massive constituent."
    },
    "1402/1402.5362_arXiv.txt": {
        "abstract": "Until recently, the detection of magnetic fields at the surface of intermediate-mass main-sequence stars has been limited to Ap/Bp stars, a class of chemically peculiar stars. This class represents no more than 5-10\\% of the stars in this mass range. This small fraction is not explained by the fossil field paradigm that describes the Ap/Bp type magnetism as a remnant of an early phase of the star-life. Also, the limitation of the field measurements to a small and special group of stars is obviously a problem to study the effect of the magnetic fields on the stellar evolution of a typical intermediate-mass star.   Thanks to the improved sensitivity of a new generation of spectropolarimeters, a lower bound to the magnetic fields of Ap/Bp stars, a two orders of magnitude desert in the longitudinal magnetic field and a new type of sub-gauss magnetism first discovered on Vega have been identified. These advances provide new clues to understand the origin of intermediate-mass magnetism as well as its influence on stellar evolution. In particular, a scenario has been proposed whereby the magnetic dichotomy between Ap/Bp and Vega-like magnetism originate from the bifurcation between stable and unstable large scale magnetic configurations in differentially rotating stars. In this paper, we review these recent observational findings and discuss this scenario. ",
        "introduction": "The origin of the stellar magnetic fields and their effects on the star structure and evolution are the two basics questions of stellar magnetism. In all stars that possess a convective envelope, a magnetic field is believed to be generated by dynamo mechanism. It is not always possible to detect it, but its presence at the surface of these solar-type stars makes little doubt. For these stars, we also know that a magnetized wind exerts a breaking torque that strongly affects their angular momentum evolution. This illustrates that although limited our level of understanding of the magnetic fields of solar-type stars enables to answer at least some simple questions about their origin and their impact. This is not the case for hotter stars with radiative envelope where the origin of the observed fields as well as their impact on the star evolution remain largely mysterious. The properties of the observed fields are nevertheless well established. They are large scale mostly dipolar fields with dipole strength ranging from 300 G to 30 kG and are remarkably stable over time (Donati \\& Landstreet, 2009). These fields are usually called fossil fields as a reference to the hypothesis that describes them as remnant of an early phase of the star-life, either from the collapse of a magnetized cloud or during the convective protostellar phase. But they only are detected in a small 5-10\\% fraction of the intermediate-mass and massive stars and this fact did not receive a convincing explanation yet. From the point of view of stellar evolution, such a lack of understanding of the field generation processes combined with the absence of direct measurements for the vast majority (90-95 \\%) of stars are a major obstacle to model the effect of a magnetic field on the structure and evolution of typical intermediate-mass and massive stars.    In this paper, we review recent observational results that shed a new light on the 5-10\\% problem and suggested new scenarios for the origin of hot star magnetism. These results have been obtained thanks to the high sensitivity of a new generation of high-resolution \\'echelle spectropolarimeters (MuSiCoS, Narval, ESPaDOnS, HARPSpol). In particular, their large wavelength coverage allows to sum up the magnetically polarized signal of many spectral lines and thus to detect weak fields. The number of lines available being an important factor, it is comparatively easier to detect weak fields in cooler stars. For example, \\cite{AK09} showed that a sub-gauss longitudinal field can be detected on the bright late-type star Pollux averaging only a few high S/N Stokes V of Narval or ESPaDOnS. For the same reason, weak fields are easier to detect in intermediate-mass stars than in more massive stars.  Spectropolarimetric surveys of Ap stars as well as of A-type non-Ap stars have first ruled out the possibility that the 5-10\\% fraction simply reflects the detection limit of magnetic measurements. Indeed, the magnetic fields of Ap/Bp stars were found to be higher than a lower limit of the order of 100 G for the longitudinal field while no field were detected in A-type non-Ap stars. Then two detections of sub-gauss fields on the bright A stars Vega and Sirius have been obtained. Thus, instead of a single class of magnetic stars, intermediate-mass star magnetism is now characterized by two type of magnetisms, Ap/Bp and Vega-like, separated by a two orders of magnitude magnetic desert between 1G - 100 G. In section 2, we present the evidences for the lower bound to the Ap/Bp magnetic fields and the magnetic desert, then the discovery of a new type of sub-gauss magnetism is described in section 3. The implications for the origin of upper-main-sequence magnetism are discussed in section 4. ",
        "conclusions": "While some years ago the fraction of intermediate-mass stars hosting large-scale stable magnetic fields like Ap/Bp stars was unknown, this class of magnetic stars is now known to be clearly separated from the other stars by a two order of magnitude magnetic desert in longitudinal field down to a new type of sub-gauss magnetism. This paradigm shift opens new perspectives for both observational and theoretical studies. Investigating how the lower bound of Ap/Bp magnetic fields depends on the stellar parameters could help to discriminate between the different hypothesis regarding the origin of this magnetism. For example we have seen that the Auri\\`ere et al. scenario predicts that this lower bound is proportional to the rotation rate. Moreover, the nature and the incidence of Vega-like magnetism among non Ap/Bp stars need to be elucidated although detecting many sub-gauss fields in A stars will be very demanding in observing time. Another important question is whether the dichotomy of intermediate-mass star magnetism is also present among pre-main-sequence intermediate-mass stars and massive stars. This is already suggested by recent spectropolarimetric surveys that detected fossil-like magnetic fields with an incidence similar to Ap/Bp stars (Alecian et al. 2013, Wade et al. 2013).  Various scenarios have been proposed to account for the magnetic dichotomy. Here we described the bifurcation between stable and unstable magnetic configurations and briefly mentioned the merging scenario for Ap/Bp stars and the failed fossil scenario for Vega-like magnetism (see the contributions by N. Langer and J. Braithwaite in these proceedings for more details). Besides the confrontation with observations, numerical simulations will also help to test these ideas. MHD computations in polytropic stars already enabled to find stable magnetic configurations that might correspond to Ap/Bp stars (Braithwaite 2009). But, except for very strong fields, differential rotation should also play an important role in shaping magnetic configurations in radiative interiors and this effect needs to be fully explored numerically."
    },
    "1402/1402.1631_arXiv.txt": {
        "abstract": "% The Pan-STARRS1 (\\PS) survey has obtained imaging in 5 bands (\\grizy) over 10 Medium Deep Survey (MDS) fields covering a total of 70 square degrees. This paper describes the search for apparently hostless supernovae (SNe) within the first year of PS1 MDS data with an aim of discovering  superluminous supernovae (SLSNe). A total of 249 hostless transients were discovered down to a limiting magnitude of $M_{\\rm AB}\\sim23.5$, of which 76 were classified as type Ia SNe. There were 57 SNe with complete light curves that are likely core-collapse SNe (CCSNe) or type Ic SLSNe and 12 of these have had spectra taken. Of these 12 hostless, non-type Ia SNe, 7 were SLSNe of type Ic at redshifts between 0.5-1.4. This illustrates that the discovery rate of type Ic SLSNe  can be maximised by concentrating on hostless transients and removing normal SNe\\,Ia. We  present data for two possible SLSNe; PS1-10pm ($z=1.206$) and PS1-10ahf ($z=1.1$), and estimate the rate of type Ic SLSNe to be between $3^{+3}_{-2}\\times10^{-5}$ and $8^{+2}_{-1}\\times10^{-5}$ that of the CCSN rate within $0.3\\leq z\\leq1.4$ by applying a Monte-Carlo technique. The rate of slowly evolving, type Ic SLSNe (such as SN2007bi) is estimated as a factor of 10 lower than this range. ",
        "introduction": "\\label{intro}  Optical transients which are spatially coincident or associated with elongated and extended sources have a high probability of being supernovae (SNe).  Thus when searching for SNe, concentrating on high mass or intrinsically high luminosity galaxies can be a fruitful endeavour to optimise the yield of recorded events as it will maximise the number of stars observed that can potentially explode as SNe.  To date the majority of such searches at low redshift have adopted this approach to good effect, for example the Lick Observatory Supernova Search \\citep{loss}.  The unbiased nature of  surveys like the Palomar Transient Factory \\citep[PTF,][]{dwarfrates} and Pan-STARRS1 \\citep[PS1,][]{PS1_system} mean that transients are being discovered without a bright galaxy bias and the neighbourhoods in which SNe are being found are not restricted to large, star forming galaxies.  \\cite{dwarfrates} found that the core-collapse SNe (CCSNe) population in dwarf galaxies is different to that found in giant galaxies, in the sense that there are many more broad-lined, Type Ic SNe in the former population. They link this to the metallicity of the underlying stellar population and its effect on stellar evolution.  Another class of SNe that have so far been discovered almost exclusively hosted by smaller, fainter galaxies is the relatively rare breed of superluminous SNe (SLSNe). \\cite{bluedeath} unravelled the mysteries of the luminous SN2005ap and SCP 06F6 by grouping them with a number of PTF discoveries to suggest these SLSNe as the death throes of at least some of the most massive of stars. Detailed studies of some of these events, such as SN2010gx \\citep[also known as PTF10cwr,][]{10gx,bluedeath} and PTF12dam \\citep{12dam}, have increased the knowledge base on these highly energetic events and the recent \\PS\\ discoveries of  PS1-10ky, PS1-10awh \\citep{10kyawh},  PS1-10bzj \\citep{lun} and PS1-11ap \\citep{11ap} are supporting and expanding progenitor and physical explosion mechanism theories. The highest redshift discoveries ($z>1.5$) from PS1 of \\cite{berger} and from the Supernova Legacy Survey \\citep[SNLS,][]{2013arXiv1310.0470H} illustrate that SLSNe can be spectroscopically followed at significantly higher redshifts than Type Ia SNe (SNe\\,Ia).  \\cite{bluedeath} used the identification of narrow Mg\\,{\\sc ii} $\\lambda\\lambda$2796,2803 absorption lines from foreground gas to place robust lower limits on the redshift of these SLSNe, which immediately provided an estimate of the enormous luminosities. In a few cases the redshifts of the Mg\\,{\\sc ii} absorption exactly matches the emission lines of the host galaxy, confirming the reasonable assumption that the Mg\\,{\\sc ii} absorption arises in the host galaxy itself. The redshifts derived by \\cite{bluedeath} and then by \\cite{10kyawh} using this method find peak absolute SLSNe magnitudes of $M_{u}\\simeq-22\\pm0.5$\\,mag and total radiated energies $\\gtrsim10^{51}$\\,erg, making them substantially more luminous than any other SN-type events.   \\begin{table*} \\caption{\\PS\\ Medium Deep Field Centres. } \\label{table:fields} \\begin{tabular}{lrr} \\hline \\hline {\\bf Field} & {\\bf RA (degrees, J2000)} & {\\bf Dec (degrees, J2000)} \\\\ \\hline MD00  &  10.675 & $ 41.267$ \\\\ MD01  &  35.875 & $ -4.250$ \\\\ MD02  &  53.100 & $-27.800$ \\\\ MD03  & 130.592 & $ 44.317$ \\\\ MD04  & 150.000 & $  2.200$ \\\\ MD05  & 161.917 & $ 58.083$ \\\\ MD06  & 185.000 & $ 47.117$ \\\\ MD07  & 213.704 & $ 53.083$ \\\\ MD08  & 242.787 & $ 54.950$ \\\\ MD09  & 334.188 & $  0.283$ \\\\ MD10  & 352.312 & $ -0.433$ \\\\ MD11  & 270.000 & $ 66.561$ \\\\ \\hline \\end{tabular} \\medskip \\end{table*}   \\begin{table*} \\caption{\\PS\\ Medium Deep Survey, typical cadence. FM$\\pm$3 designates 3 nights on either side of Full Moon.} \\label{table:cadence} \\begin{tabular}{ccc} \\hline \\hline {\\bf Night} & {\\bf Filter} & {\\bf Exposure Time} \\\\ \\hline 1         & \\gps \\& \\rps & 8$\\times$113s each\\\\ 2         & \\ips         & 8$\\times$240s \\\\ 3         & \\zps         & 8$\\times$240s \\\\ repeats... & . . . & . . . \\\\ FM$\\pm3$  & \\yps         & 8$\\times$240s \\\\ \\hline \\end{tabular} \\medskip \\end{table*}  Despite this common lower limit, differences in the photometric and spectroscopic evolution of the observed SLSNe suggests a number of progenitor possibilities. \\cite{10gx} find iron and other features normally associated with Type Ic SNe in the spectra of \\gx\\ at $30-50$ days after peak showing that the transient evolved to resemble an energetic Type Ic SN but on a much slower timescale.  \\cite{11xk} present data on another five such objects at redshifts 0.1 - 0.2 with detailed modelling suggesting that the explosions are simply `normal' Type Ic SNe with an additional power source providing a boost to the luminosity.  The best fitting models presented by \\cite{11xk} are those where a fast rotating neutron star (magnetar) provides the extra energy required, an idea proposed and developed by \\cite{pulsara}, \\cite{pulsarb} and \\cite{magn1}.  Other power sources, such as the radioactive decay of $^{56}$Ni \\citep{arnett}, are also explored but could not be reconciled with the light curves of SLSNe. These kind of models were used by \\cite{12dam} and \\cite{11ap} to explain the slower photometric evolution of the SLSNe \\dam\\ and \\ap, where again the magnetar models gave the most satisfactory fits to the data.  Prior to the discovery of these two SLSNe however, the only well studied object of this slowly evolving type was \\bi\\ \\citep{07bi,DY} which was believed to have been the result of a pair-instability supernova \\citep[PISN,][]{1stpisn, 2ndpisn, 3rdpisn, 4thpisn}. \\cite{snsh} suggest the interaction of the SN shock with a dense circumstellar material of H-poor material as a possible mechanism for the production of the observed features of the light curves and spectra.  This theory also accounts for the Ic-like features and the unusually high magnitudes associated with the aforementioned objects. Thus we will refer to these two subclasses of SLSNe, those sharing the properties of the \\cite{bluedeath} sample and those that display the prolonged light curve evolution of SN2007bi, as SLSNe-Ic and slowly evolving SLSNe-Ic respectively.  A number of discoveries of a completely separate class of SLSNe that display strong H emission similar to the Type II classes of SNe have also been made.  SN2006gy \\citep{06gya,06gyb} and SN2003ma \\citep{03ma} are examples of this class, where the observed features are likely produced by the interaction between an  energetic SN explosion and a very dense circumstellar medium. Any objects of this class will be referred to here as SLSNe-II \\citep{gal-yam}. In contrast to the SLSNe-Ic,  SN2006gy and SN2003ma occured in bright host galaxies, suggesting that SLSNe-II do not follow the apparent trend visible for SLSNe Ic which have, almost exclusively, dwarf hosts.  An important feature of the known SLSNe-Ic sample is the noticeable preference for them appearing in faint host galaxies. This trend has been apparent right from the early discoveries with a faint host galaxy of no brighter than $M_{r}\\sim-18$  found for SN2005ap \\citep{05ap} and an upper  limit of $-18.1$ mag set at the pre-explosion location of SCP06F6 \\citep{scp06f6}.  \\cite{10gxgal} and \\cite{bluedeath} found a faint dwarf host for SN2010gx in archive SDSS images (SDSS J112546.72-084942.0), with an absolute magnitude of $M_r\\sim-18$ and a more recent study  further refined this value to $M_g=-17.42\\pm0.17$ \\citep{janet}. \\cite{DY} also calculated an absolute magnitude for the host of SN2007bi using SDSS archive images (SDSS J131920.14+085543.7) and found an $M_{B}\\sim-16.4$. A similar trend has been noted with virtually all subsequent discoveries \\citep{bluedeath,10kyawh,10gx,06oz,scp06f6, janet,lun,2014arXiv1405.1325N}, which led \\cite{2013arXiv1311.0026L} to study a large sample of host galaxies and suggest that they are similar hosts to those of Gamma-Ray Bursts (GRBs).  An exception to this supposed trend is the event so far reported with one of the highest redshifts. PS1-11bam \\citep{berger}, which has $z=1.55$ derived from strong absorption of both Fe\\,{\\sc ii} and Mg\\,{\\sc ii} and the detection of [OII] 3727\\AA\\ from the host galaxy in emission, has the most luminous host discovered so far for a SLSN-Ic, with a near UV absolute magnitude of M$_{\\rm NUV} \\simeq -20.3$. However the host is still some 2 magnitudes fainter than the SN itself.  In summary, the currently known SLSNe-Ic sample typically appear to be $>2-4$ mag brighter than their hosts \\citep[see][for a compilaton of host magnitudes]{2013arXiv1311.0026L} indicating that a simple way of isolating them in higher redshift searches could be to target transients with either no host detected or with a significant difference between total host luminosity and peak SLSN-Ic magnitude. It is striking that no SLSNe-Ic erupting in bright galaxies have  been uncovered by any previous wide-field survey or galaxy targeted low redshift searches. Although one might argue there is always a preference for observers to take spectra of isolated transients to avoid galaxy contamination effects, the very large sample of low redshift supernovae now classified do not contain any examples of SLSNe Ic in galaxies close to $L^{\\star}$  \\citep[the characteristic luminosity in the Schechter luminosity function for galaxies, approximately corresponding to $M_{B}\\simeq-21$][]{1976ApJ...203..297S}. This is somewhat surprising since the bulk of stellar mass, and star formation, occurs in galaxies close to $L^{\\star}$.  The focus of this paper is to use this apparent preference for dwarf galaxy hosts as a  method of finding SLSNe-Ic in the first year of the \\PS\\ Medium Deep Survey (\\MDS),  to quantify their numbers in a magnitude limited survey and to approximately estimate the volumetric rates between the redshift range of $z = 0.3 - 1.4$. \\cite{quimbrate} previously estimated the SLSN-Ic rate to be $32^{+77}_{-26}$ events Gpc$^{-3}$yr$^{-1}h^{3}_{71}$ at a redshift of $z\\sim0.2$ (although based on 1 event) and the SLSN-II rate to be $151^{+151}_{-82}$ events Gpc$^{-3}$yr$^{-1}h^{3}_{71}$ at $z \\sim 0.15$.   At the beginning of the  PS1 survey we used reference images which were made of a small number  ($\\sim8$) of single input images from a good night. As the survey has progressed we have been able to stack the images to build much deeper template and stack images. However at the beginning of the PS1 survey and during the first year of operations, we used reference stacks which reached a limiting apparent magnitudes of $\\sim23.5$ in the \\gps\\rps\\ips-bands. Hence SNe which are brighter than 22, but have no host visible brighter than 23.5 were immediately candidates for hunting down these exciting phenomena. This paper focuses on all hostless transients discovered during the first year of the \\PS\\ survey operations. ",
        "conclusions": "\\label{sec:conclusion}  We have catalogued all the SN-like, hostless transients from the PS1 Medium Deep Survey and, by using multiple spectroscopic programmes and photometric classifiers, filtered out all the SNe\\,Ia events. This leaves a promising percentage of remaining objects that seem to fall into the category of SLSNe-Ic.  \\begin{itemize} \\item 249 hostless transients were discovered within the first 1.37 yr of the \\PS\\ \\MDS, 133 of which have SN-like features. \\item 40 are spectroscopically confirmed SNe, $\\sim17.5\\%$ of which are possible SLSNe-Ic. \\item 12 are spectroscopically confirmed, non-type Ia SNe, $\\sim60\\%$ of which are possible SLSNe-Ic. \\end{itemize}  PS1-10pm and PS1-10ahf were discovered in this way. Photometric and spectroscopic comparisons place PS1-10pm comfortably in the SLSNe-Ic class.  The classification of PS1-10ahf is not as robust, but reasonably solid photometric and spectroscopic comparisons give it a probable association with the slowly evolving class of SLSNe-Ic such as SN2007bi, PS1-11ap and PTF12dam. We highlight that this was a combination of spectroscopic classification (when the SNe were close to peak) and photometric classification after the lightcurves had been gathered. A challenge remains to carry out accurate photometric classification in real time.  Using the SLSNe-Ic statistics gathered during the search for orphans and comparing them with Monte-Carlo simulations of SLSNe-Ic, we determined the rate of SLSNe-Ic within a redshift of $0.3\\leq z\\leq1.4$ to be between  $3^{+3}_{-2}\\times10^{-5}$ and $8^{+2}_{-1}\\times10^{-5}$ that of the CCSNe rate. The ratio of slowly evolving SLSNe-Ic to CCSNe seems to be much lower, at around $3^{+4}_{-2}\\times10^{-6}$.  Using a combination of careful photometric analysis and thorough spectroscopic follow-up and the search method of exploiting the common characteristic of the $>2$ mag difference between discovered SLSNe-Ic peak magnitudes and their host galaxies, an ever increasing number of SLSNe-Ic should be found in the next few years from current and future wide-field surveys (Pan-STARRS2, The Dark Energy Survey, The Zwicky Transient Factory, La Silla-QUEST + the Public ESO Spectroscopic Survey of Transient Objects, Large Synoptic Survey Telescope). \\\\ \\\\ {\\it Facilities:} Pan-STARRS1, Gemini, William Herschel Telescope  {\\small{ \\textit{Acknowledgements}.  The Pan-STARRS1 Surveys (PS1) have been made possible through contributions of the Institute for Astronomy, the University of Hawaii, the Pan- STARRS Project Office, the Max-Planck Society and its participating institutes, the Max Planck Institute for Astronomy, Heidelberg and the Max Planck Institute for Extraterrestrial Physics, Garching, The Johns Hopkins University, Durham University, the University of Edinburgh, Queen's University Belfast, the Harvard- Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global Telescope Network Incorporated, the National Central University of Taiwan, the Space Telescope Science Institute, the National Aeronautics and Space Administration under Grant No. NNX08AR22G issued through the Planetary Science Division of the NASA Science Mission Directorate, the National Science Foundation under Grant No. AST-1238877, and the University of Maryland. S.J.S. acknowledges funding from the European Research Council under the European Union's Seventh Framework Programme (FP7/2007-2013)/ERC Grant agreement no [291222] (PI: S.J. Smartt). This work is based on observations made with the following telescopes: William Herschel Telescope (operated by the Isaac Netwon Group), in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofi\u0301sica de Canarias, in the island of La Palma; the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Minist\\'{e}rio da Ci\\^{e}ncia, Tecnologia e Inova\\c{c}\\~{a}o (Brazil) and Ministerio de Ciencia, Tecnolog\\'{i}a e Innovaci\\'{o}n Productiva (Argentina). Some observations reported here were obtained at the MMT Observatory, a joint facility of the Smithsonian Institution and the University of Arizona.  Support for SR was provided by NASA through Hubble Fellowship grant \\#HST-HF-51312.01 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.  JT acknowledges support for this work provided by National Science Foundation grant AST-1009749.  SM acknowledges financial support from the Academy of Finland (project: 8120503). We thank A. Gal-Yam and P. Nugent for providing the classification of PS1-11acn, which is the PTF object PTF11dws.}"
    },
    "1402/1402.4687_arXiv.txt": {
        "abstract": "When muons travel through matter, their energy losses lead to nuclear breakup (``spallation'') processes.  The delayed decays of unstable daughter nuclei produced by cosmic-ray muons are important backgrounds for low-energy astrophysical neutrino experiments, e.g., those seeking to detect solar neutrino or diffuse supernova neutrino background (DSNB) signals.  Even though Super-Kamiokande has strong general cuts to reduce these spallation-induced backgrounds, the remaining rate before additional cuts for specific signals is much larger than the signal rates for kinetic energies of about 6 -- 18 MeV.  Surprisingly, there is no published calculation of the production and properties of these backgrounds in water, though there are such studies for scintillator.  Using the simulation code FLUKA and theoretical insights, we detail how muons lose energy in water, produce secondary particles, how and where these secondaries produce isotopes, and the properties of the backgrounds from their decays.  We reproduce Super-Kamiokande measurements of the total background to within a factor of 2, which is good given that the isotope yields vary by orders of magnitude and that some details of the experiment are unknown to us at this level.  Our results break aggregate data into component isotopes, reveal their separate production mechanisms, and preserve correlations between them.  We outline how to implement more effective background rejection techniques using this information.  Reducing backgrounds in solar and DSNB studies by even a factor of a few could help lead to important new discoveries. ",
        "introduction": "\\label{sec:intro} Neutrinos are powerful probes of the universe and its contents.  They are abundantly produced by nuclear fusion processes that convert protons into neutrons, through the decays of unstable particles and nuclei created in high-energy processes, and through pair production in hot, dense environments.  They can reach us unattenuated and undeflected from vast distances or from behind enormous column densities of matter, directly revealing the energies and timescales of the processes that made them.  Even in a core-collapse supernova, where the neutrinos are thermalized by scattering, they emerge at energies $\\sim 10$ MeV over about 10 s, compared to photons, which emerge at energies $\\sim 1$ eV over months.  The  detection of astrophysical neutrinos allows us to probe physical conditions and neutrino properties beyond the reach of laboratory experiments.  The first great challenge of neutrino astronomy is the fact that the small interaction cross sections that make the above possible make detection difficult.  This can only be solved by brute force --- building large enough detectors to ensure adequate event rates.  We focus on Super-Kamiokande (Super-K), the world's largest low-energy neutrino detector, which has a fiducial mass of 22.5 kton of water and a total mass of 50 kton of water~\\cite{Fukuda2003,Abe2014}.  (For comparison, neutrinos were first detected in the Reines-Cowan reactor experiment with a detector using less than 1 ton of scintillator~\\cite{Cowan1956}.)  Even with such a large detector, the measured rates of low-energy astrophysical neutrinos are very small: about 15 solar neutrino events (all flavors of neutrinos elastically scattering electrons) detected per day~\\cite{Hosaka2006,Cravens2008,Abe2011} and an upper limit of several events (primarily $\\bar{\\nu}_e$ inverse beta decay) detected per year from the diffuse supernova neutrino background (DSNB)~\\cite{Malek2003,Horiuchi2009,Beacom2010,Bays2012,Zhang2013}.  The second and far greater challenge of neutrino astronomy is reducing detector backgrounds to isolate these rare signals.  Immense care and sophistication is required, and continual progress with existing detectors is possible.  The primary backgrounds for solar and DSNB signals are MeV electrons and positrons from the decays of nuclei and muons.  Below about 6 MeV detected electron kinetic energy, intrinsic radioactivities are the dominant background in Super-K~\\cite{Gando2003,Hosaka2006,Cravens2008,Abe2011}, and these are controlled through the selection and purification of materials, choice of water circulation pattern to minimize radon ingress, and software processing (e.g., reconstruction quality and fiducial volume cuts).  From about 6 to 18 MeV kinetic energy, induced radioactivities produced by cosmic-ray muons are the dominant background~\\cite{Gando2003,Hosaka2006,Cravens2008,Abe2011}, and there is great potential to reduce these with the help of theoretical work.  To reduce cosmic-ray backgrounds, Super-K was built under 1000 m of rock (2700 m water equivalent) in the Kamioka mine in Japan~\\cite{Fukuda2003,Abe2014}.  As cosmic-ray particles interact with the rock and lose energy, their flux is reduced. The only high-energy particles that reach the Super-K detector are muons and neutrinos.  The muon flux is $6.0 \\times 10^5$ m$^{-2}$ hr$^{-1}$ at sea level, and is reduced to 9.6 m$^{-2}$ hr$^{-1}$ at Super-K~\\cite{Galbiati2005}, which corresponds to a muon rate in the detector of about 2 Hz~\\cite{Hosaka2006}.  It is easy to veto the muons themselves, but they frequently produce relatively long-lived radioactive isotopes through the breakup (``spallation\") of stable nuclei directly or, more commonly, through secondary particles produced through muon energy-loss processes.  The spallation rate is large, $\\sim 1$ interaction per through-going muon in Super-K, though many of the daughter nuclei are stable or decay in ways that do not produce Cherenkov signals.  Super-K has cuts to reduce backgrounds from the decays of spallation products, but these have to be limited to not overly discard signal events.  Many of the unstable isotopes produced have half-lives of order 1 s, comparable to the time between successive muons.  It is easy to estimate that a simple cut of all events in a cylinder of radius even a few meters around each muon track for a few seconds leads to a detector deadtime of $\\sim 20\\%$.  The real algorithm used by Super-K is more complex, and is based on a likelihood analysis that takes into account distance and time from the preceding muon as well as a variable related to muon energy loss, but a similar deadtime is achieved~\\cite{Koshio, Hosaka2006}.  Even though the Super-K spallation cuts have a rejection efficiency of $\\sim 90\\%$~\\cite{Gando2003}, the remaining background rate is still $\\sim 10$ times greater than the solar neutrino signal rate above several MeV (this is then reduced by another factor $\\sim$ 10 by the solar direction cut, leaving a background comparable to the signal)~\\cite{Abe2011}.  For the DSNB search, a higher energy threshold can be used to dramatically reduce backgrounds, but spallation decays are still overwhelming below about 18 MeV~\\cite{Malek2003} (16 MeV with new techniques~\\cite{Bays2012}).  Our goal for this paper is to detail the production processes for spallation backgrounds in Super-K and the physical characteristics of where, when, and with what associated particles these decays occur.  With this information, it will be possible to make better cuts to reject backgrounds while preserving signals.  For solar neutrinos, such improvements could help improve the significance of the 2.7-$\\sigma$ hint of the day-night effect from neutrino mixing in Earth~\\cite{Liu1997,Maris1997,Blennow2004,Renshaw2013}.  They may also help lead to the first detection of the {\\it hep} neutrino flux, which is likely only a factor of a few away from detection~\\cite{Marcucci2000,Bahcall2001,Fukuda2001,Bahcall2004,Hosaka2006}.  Such measurements would improve our knowledge of the Sun and of neutrino mixing parameters~\\cite{Fogli2011,Haxton2012,Bellini2013lSolar}.  Reduction of spallation backgrounds would also help lower the energy threshold in the DSNB search~\\cite{Malek2003,Bays2012} to where the signal is larger~\\cite{Horiuchi2009,Beacom2010}, which might help lead to a first detection.  Until now, there has been no detailed published study of spallation backgrounds in water.  The Super-K cuts have been developed from empirical studies~\\cite{Hosaka2006,Cravens2008,Abe2011, Bays2012}, and not from theoretical calculations.  Further, they treat all isotopes together, without taking into account significant differences in their production, properties, and distributions.  With Super-K nearly reaching the sensitivity needed for the above discoveries, a more detailed approach is needed. The interactions of muons with scintillator have been studied extensively with underground~\\cite{Aglietta1989, Hertenberger1995, Boehm2000, Menghetti2005} and accelerator~\\cite{Hagner2000,Chazal2002,Lin2013} experiments, and measurements like these have been incorporated into simulation packages like FLUKA~\\cite{Ferrari2005,Battistoni2007} and GEANT4~\\cite{Agostinelli2002,Allison2006}.  This gives an opportunity to check our work and to understand the expected uncertainties of the simulations.  This paper is not meant to be a comprehensive study.  It is a first step in understanding spallation backgrounds in water-based detectors, beginning with the yields and the average physical distributions of secondaries and isotopes.  In two subsequent papers, we will go further, showing how characteristics of the showers of secondary particles that produce isotopes can be used to tailor better cuts~\\cite{Bays2012} and how those would be improved if Super-K gained the ability to detect neutrons by adding dissolved gadolinium~\\cite{Beacom2004}.  This paper is organized as follows.  In Sec.~\\ref{sec:setup}, we describe the setup for our simulation.  In Sec.~\\ref{sec:energy_loss}, general points about muon energy loss and secondary particle production are discussed. Our main results are in Sec.~\\ref{sec:backgrounds}, where we calculate the neutron and isotope yields and study the properties of the induced backgrounds.  Finally, we present our conclusions in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Guided by theoretical understanding and analysis, we use the simulation package FLUKA to study muon interactions with water, the production and properties of secondary particles, and the production and decay of unstable isotopes.  Where possible, we compare our results to published measurements from Super-K, finding good agreement on an absolute scale, i.e., a factor of 2, which is reasonable considering the orders of magnitude differences in production rates.  The residual discrepancies primarily arise from uncertainties in hadronic interactions and unpublished details of the muon backgrounds, and some of the differences could be reduced by calibration to measured data.  As a check, we also performed similar calculations for scintillator-based detectors, for which there are more extensive theoretical studies and experimental measurements.  We focus on comparison to isotope and neutron production in KamLAND~\\cite{Abe2010} and Borexino~\\cite{Bellini2013}, finding good agreement, within the factors of 2 that have been noted by others between the measurements and calculations and also between calculations with FLUKA versus GEANT4~\\cite{Wang2001, Kudryavtsev2003}.  One interesting point for context is how different the spallation backgrounds are in water Cherenkov detectors compared to scintillator detectors.  First, for water there is the fortunate point that, although the production rate of all isotopes is comparable to that in scintillator, that of unstable isotopes is about ten times less.  Second, many of the unstable isotopes decay without producing Cherenkov light.  Scintillator detectors have the ability to detect neutrons through their radiative captures, and this is a significant advantage in identifying spallation products.  However, if Super-K adds dissolved gadolinium to enable the detection of neutron captures, it will have a similar capability~\\cite{Beacom2004}.  Our calculations for Super-K lead to important new high-level results beyond the details presented here.  First, a demonstration that a theoretical calculation of the spallation backgrounds in water is now possible, even though it was not when Super-K began~\\cite{Koshio}.  Compared to an empirical approach, production mechanisms are revealed, aggregates are separated into components, and correlations are preserved.  Second, we show details that were heretofore unavailable.  Important examples are differences between the distributions and correlations of each isotope, including temporal distribution after the muon, distance distribution away from the muon, decay energy spectrum, and associated particles.  We demonstrate that there is more information to be gained by having likelihood functions of time and distance for each isotope.  Instead of a global likelihood for all spallation decays, our results could be used to construct per-isotope likelihoods that would lead to more precise cuts.  Also, a new variable of decay energy can be used in addition to its original three variables of decay distance to the muon track, decay time, and muon energy loss.  Even modest improvements, say a factor of a few, could lead to significant gains in the ability to measure signals. This could help lead to first discoveries of the day-night effect and the {\\it hep} flux in solar neutrinos, as well as the DSNB.  Our results are calculated for Super-K, but they could have wider applicability.  The isotope yields per muon vary only moderately with depth, once that depth is appreciable, because they have a modest dependence on the muon average energy, scaling roughly as $E_\\mu^{0.8-1.1}$~\\cite{Abe2010}.  As first estimates, our results would provide useful comparisons for the Sudbury Neutrino Observatory~\\cite{Aharmim2013}, Hyper-Kamiokande~\\cite{Abe2011hk}, and the water shields of a variety of neutrino and dark matter detectors.  It would be valuable for Super-K to produce a dedicated study on spallation backgrounds informed by the predictions of this paper.  The yields of different isotopes could be identified by a global fit that takes into account the full energy and time information on spallation decays, e.g., energy spectra in different time ranges, as has been done for scintillator detectors~\\cite{Abe2010,Bellini2013}.  Another key observable is the radial distributions of isotopes produced by different types of secondaries.  An improved FLUKA simulation could be developed using a more complete description of the detector details, especially the muon distributions.  It seems likely that the uncertainties could be reduced to well below a factor of 2 by calibrating the simulation to measured data.  It would also be valuable to have a similar study for the Sudbury Neutrino Observatory~\\cite{Aharmim2013}.  The very low muon rate and intrinsic radioactivities would make it easier to identify spallation decays and to avoid confusion over which muon was the parent.  In addition, the ability to detect neutrons would help identify isotope production channels.  With corrections for the different muon spectrum, detector properties, and the production of neutrons by deuterium photo-disintegration, it would be straightforward to relate these measurements to Super-K results.  In two follow-up papers, we will develop further ways to reduce backgrounds in Super-K and other water Cherenkov detectors.  In the first paper, we will study the variations in muon energy loss along the path due to showers, and how this can be used to identify where isotopes are produced.  This effect was discovered empirically in Ref.~\\cite{Bays2012}, and our results will provide the first detailed explanation of how it works and how it could be improved.  In the second paper, we will show how the ability to detect neutrons using gadolinium in water, as first suggested in Ref.~\\cite{Beacom2004}, can be used to improve cuts to reduce spallation backgrounds.  These papers will include some surprises that will allow significant gains in sensitivity beyond those enabled by results given here."
    },
    "1402/1402.3772_arXiv.txt": {
        "abstract": "We report on the results of high dispersion spectroscopy of two `superflare stars', KIC 9766237, and KIC 9944137 with Subaru/HDS. Superflare stars are G-type main sequence stars, but show gigantic flares compared to the Sun, which have been recently discovered in the data obtained with the Kepler spacecraft.  Though most of these stars are thought to have a rotation period shorter than 10 days on the basis of photometric variabilities, the two targets of the present paper are estimated to have a rotation period of 21.8 d, and 25.3 d.  Our spectroscopic results clarified that these stars have stellar parameters similar to those of the Sun in terms of the effective temperature, surface gravity, and metallicity.  The projected rotational velocities derived by us are consistent with the photometric rotation period, indicating a fairy high inclination angle.  The average strength of the magnetic field on the surface of these stars are estimated to be 1-20 G, by using the absorption line of Ca \\textsc{ii} 8542.  We could not detect any hint of binarity in our spectra, although more data are needed to firmly rule out the presence of an unseen low-mass companion.  These results claim that the spectroscopic properties of these superflare stars are very close to those of the Sun, and support the hypothesis that the Sun might cause a superflare. ",
        "introduction": "A flare event on the Sun was visually noticed, and recorded by \\citet{car1859} for the first time.  Since then, many solar flares have been observed in all the wavelengths, and these are now generally explained by explosive release of the magnetic energy stored around sunspots (for a review, e.g. \\cite{shi11}). The total energy released in the largest solar flares was estimated to be of the order of 10$^{32}$ erg (e.g. \\cite{pri81, ems12}), while a possibility has been recently discussed that the most energetic solar flares caused extreme cosmic-ray events evidenced by isotopic abundance research of historical relics \\citep{miy12, uso12, mel12, miy13}.  Such strong solar flares could have a severe effect on the terrestrial environment.  Many `superflares' were discovered in G-type main sequence stars, namely solar-type stars, in 2012, and 2013 \\citep{mae12, shi13b}, while only nine such events had been reported before it \\citep{sch00}.\\footnote{Some of these nine flares seem doubtful. Gmb 1830 (=HR 4550) have a flare star companion, CF UMa (see \\cite{rei91}).  The superflare visually observed on S For is suggested to be due to mis-identification (see \\cite{pay52, ash59}). In addition, $\\kappa$ Cet and $\\pi^1$ UMa rotate with a period of 9.24 d, and 4.89 d, respectively, and are rather young (age $<$ 1 Gy) \\citep{lin12}.}  Superflares are eruptive events having an energy 10 to 10$^6$ times larger than that of the largest solar flare recorded so far.  Such events have been found in many T~Tau stars, RS CVn-type binaries, and dMe stars (e.g. \\cite{shi02}). Before the advent of high-precision space photometry, however, it had been difficult to detect a superflare in white light in solar-type stars, due to contrast reasons.  Indeed, even the largest solar flares could increase the total luminosity of the Sun by only 0.03\\% even at the peak of the event (e.g. \\cite{kop05}).  This situation has been recently changed by the Kepler spacecraft which observed a sky region between Cygnus and Lyra, monitoring over 160 thousand stars with a cadence of about 30 minutes, and a quite high accuracy exceeding 0.01\\% for moderately bright stars \\citep{koc10}.  Kepler data have been then used for the stellar flare reseach on cool stars (e.g. \\cite{wal11}), and on hot stars \\citep{bal12}.  We hereafter call those solar-type stars showing superflares `superflare stars'.  The analyses of the Kepler data by \\citet{mae12} and \\citet{shi13b} revealed that the relation between the occurrence frequency of the flare and the flare energy in the Sun can be roughly extended for the flare energy up to $\\sim 10^{36}$ erg on superflare stars. \\citet{shi13b} also estimated that a superflare with an energy of 10$^{34-35}$ erg occurs once in 800-5,000 years in Sun-like stars, i.e. main-sequence stars with the effective temperature in the range 5600--6000 K, and the rotational period longer than 10 days (see also \\cite{shi13a}).  This rotation period is, however, estimated from the brightness modulation, assuming that it is due to the rotation of the star, whose photosphere is covered by large starspots (see \\cite{not13b}).  The rotation period then should be spectroscopically confirmed.  In addition, the stellar parameters, and binarity of the superflare stars should be checked for the discussion whether the Sun can really give rise to superflares, or not.  We have then performed high-dispersion spectroscopy of superflare stars in our database.  In this Letter, we report on the first results of the analysis of high-resolution spectra of two G-type stars, namely KIC 9766237 and KIC 9944137, whose rotation period was estimated to be 21.8, and 25.3 days, respectively \\citep{shi13b}. Figure \\ref{fig:lc} shows an example of the Kepler light curve of these targets, around the only one flare of each star automatically detected by the method described by \\citet{mae12}.  Quasi-periodic modulations with an amplitude of $\\sim$0.1 \\%, other than the superflare, are clearly present.  The total energy released during the superflares marked in figure \\ref{fig:lc} is estimated to be $\\sim10^{34}$ erg, and the duration is 0.1-0.2 days.  Table \\ref{tab:param} summarizes the effective temperature, surface gravity, and metallicity of KIC 9766237, and KIC 9944137, listed in the Kepler Input Catalog \\citep{bro11}.  Table \\ref{tab:param} also lists the data of 18 Sco, the best investigated solar twin (e.g. \\cite{por97, sou04}), for comparison.  \\begin{figure} \\begin{center} \\FigureFile(80mm,115mm){fig1.eps} \\end{center} \\caption{(a) Typical light curve of KIC 9766237 drawn with the long cadence Kepler data.  The Y axis represents the relative flux normalized by the average flux ($(F-F_{\\rm av})/F_{\\rm av}$). Quasi-periodic modulations with a timescale of about twenty days are seen.  A tick mark points out the superflare automatically detected by the method described in \\citet{mae12} and \\citet{shi13b}. The inset figure shows an enlarged light curve around the ticked superflare.  The amplitude is about 0.17 \\%, and the duration is about 0.1 days.  Though many `spikes' other than that superflare are seen, all of them consist of only one point, and are not judged as a superflare with our criteria.  (b) The same as (a), but of KIC  9944137.  The amplitude of the superflare in the inset figure is about 0.28 \\%, and the duration exceeds 0.2 days.} \\label{fig:lc} \\end{figure}  The details of our observation of the targets are summarized in section 2.  Section 3 describes the observational results, and brief discussion. ",
        "conclusions": "\\begin{figure} \\begin{center} \\FigureFile(80mm,115mm){Ha6563.eps} \\end{center} \\caption{H$\\alpha$ absorption line of KIC 9766237, and KIC 9944137 (solid lines), and 18 Sco (dashed lines).  The H$\\alpha$ depth of these targets is slightly smaller than that of the comparison solar-twin star, 18 Sco.  The spectrum of KIC 9944137 accompanying that of 18 Sco is shifted by +1.0 along the Y axis, for clarity.} \\label{fig:Ha} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,115mm){Ca8542.eps} \\end{center} \\caption{Absorption profile of Ca \\textsc{ii} 8542, one of Ca \\textsc{ii} IR triplet, of KIC 9766237, and KIC 9944137 (solid lines), and 18 Sco (dashed lines).  This line of both targets is slightly shallower than that of the comparison solar-twin star, 18 Sco.  The spectrum of KIC 9944137 accompanying that of 18 Sco is shifted by +1.0 along the Y axis, for clarity.} \\label{fig:Ca} \\end{figure}  Figures \\ref{fig:Ha}, and \\ref{fig:Ca} show absorption lines of H$\\alpha$, and Ca \\textsc{ii} 8542.  They are good indicators of the chromospheric activity (e.g. \\cite{fra10, tak10}).  Both lines of both targets are slightly shallower than those of the solar twin star, 18 Sco, that has the same activity level of the quiet Sun on the basis of the Mt Wilson Ca II S index (e.g. \\cite{dun91}).  The index $r_0(8542)$, the flux normalized by the continuum level at the line core of Ca \\textsc{ii} 8542 is 0.26 for KIC 9766237, 0.27 for KIC 9944137, and 0.21 for 18 Sco. These index values and effective temperature of $T_{\\rm eff} =5606 {\\rm K}$ for KIC 9766237, and $T_{\\rm eff} =5666 {\\rm K}$ for KIC 9944137 (derived later) suggest that both targets are in the region dominated by quiescent dwarf stars in figure 6 in \\citet{not13a}, though the shallowness of H$\\alpha$ and Ca II, compared with those of 18 Sco, indicates a little higher atmospheric activity level of both stars.  The average magnetic field can be roughly estimated, by using the emission core flux of Ca \\textsc{ii}.  The relation between the emission core flux of the Ca \\textsc{ii} K line and the average magnetic field was first derived for cool stars by \\citet{sch89}.  Notsu et al. (2014, in preparation) have extended this relation  to the emission core flux of the Ca \\textsc{ii} 8542 and the average magnetic field.  The average strength of the magnetic field of KIC 9766237, and KIC 9944137 is estimated to be 1-20 Gauss, by using this method.  This value is of the same order as, or one order of magnitude higher than that of the Sun ($\\sim$2 G).  An example of the photospheric absorption lines is displayed in figure \\ref{fig:Fe}.  The profile of these photospheric lines are in good agreement with that of 18 Sco.  This fact suggests that both of KIC 9766237, and KIC 9944137 are single stars, or at most single-lined binaries.  \\begin{figure} \\begin{center} \\FigureFile(80mm,115mm){Fe6212.eps} \\end{center} \\caption{Example of photospheric absorption lines, including Fe I 6213, 6215, 6216, and 6219, of the target stars, and 18 Sco, corrected for the radial velocity.  The profiles of these lines are in good agreement of that of 18 Sco, suggesting that both of the targets are single stars.  The projected rotational velocity of KIC 9766237, KIC 9944137, and 18 Sco is estimated to be 2.1($\\pm$0.3), 1.9($\\pm$0.3), and 2.2($\\pm$0.3) km s$^{-1}$, respectively.  The spectrum of KIC 9944137 accompanying that of 18 Sco is shifted by +1.0 along the Y axis, for clarity.} \\label{fig:Fe} \\end{figure}  \\begin{table*} \\caption{Stellar parameters of KIC 9766237, KIC 9944137, and 18 Sco} \\label{tab:param} \\begin{center} \\begin{tabular}{cccccccc} \\hline\\hline Name & Sp.   & $T_{\\rm eff}$ & $\\log g$ & [Fe/H] & $v \\sin i$    & A(Li) & Reference$^{\\dagger}$ \\\\ & Type  &  [K]          & [cgs]    & & [km s$^{-1}$] &  \\\\ \\hline KIC 9766237 & & 5674 & 4.557 & $-$0.147 &  & & KIC \\\\ & & 5606($\\pm$40) & 4.25($\\pm$0.11) & $-$0.16($\\pm$0.04) & 2.1($\\pm$0.3) & $<$1.0 & this study \\\\ KIC 9944137 & & 5725 & 4.619 & $-$0.064 &  & & KIC \\\\ & & 5666($\\pm$35) & 4.46($\\pm$0.09) & $-$0.10($\\pm$0.03) & 1.9($\\pm$0.3) & $<$1.0 & this study \\\\ 18 Sco & G2V & 5816 & 4.45 & 0.04 & 2.1 & 1.63 & 1,2,3 \\\\ &     & 5824($\\pm$15) & 4.50($\\pm$0.04) & 0.07($\\pm$0.02) & 2.4($\\pm$0.3) & 1.6($\\pm$0.2) & this study \\\\ \\hline \\multicolumn{8}{l}{$^{\\dagger}$1: \\citet{das12}, 2: \\citet{pet08}, 3: \\citet{mel07}}\\\\ \\end{tabular} \\end{center} \\end{table*}  An analysis of the co-added spectrum of KIC 9744237, KIC 9944137, and 18 Sco yields the stellar parameters of the effective temperature $T_{\\rm eff}$, surface gravity $\\log g$, projected rotational velocity $v \\sin i$, metallicity [Fe/H], and lithium abundance A(Li) (table \\ref{tab:param}).  The details of this analysis method were described by \\citet{not13a}.  While the values of $T_{\\rm eff}$, $\\log g$, and [Fe/H] we derived are rather different than those in the Kepler Input Catalog, spectroscopic determination of the stellar parameters are known to be more accurate than the photometric method adopted in the Kepler Input Catalog (e.g. \\cite{mol10, pin12}).  Comparison of our 18 Sco data with those in the literature supports the accuracy of our analysis.  The stellar parameters of KIC 9766237, and KIC 9944137 derived here indicate that both stars have the surface gravity and metallicity fairly close to those of the Sun, although the effective temperature of both stars is a little lower ($\\sim$100 K) than that of the Sun.  The projected rotational velocity of 2.1($\\pm$0.3) km s$^{-1}$ of KIC 9766237, and 1.9($\\pm$0.3) km s$^{-1}$ of KIC 9944137 is consistent with the photometrically estimated rotation period of 21.8, and 25.3 days, respectively, indicating that the inclination angle $i$ is fairly high, $i=90^{+0}_{-25}$ deg for KIC 9766237, and $i=75^{+15}_{-22}$ deg for KIC 9944137 (figure \\ref{fig:vsini}).  The method to derive the rotational velocity on the basis of the light curve was descibed also by \\citet{not13a}, while the data of the isochrone used here was replaced by that in \\citet{bre12}.  The spectra in the region around 6700 \\AA\\  are shown in figure \\ref{fig:Li}.  The Li I 6708 \\AA\\ line is detected neither in the spectrum of KIC 9766237 nor in that of KIC 9944137, while it is clearly visible for 18 Sco.  The Li abundance A(Li) is estimated to be smaller than A(Li)=1.0, on the basis of figure \\ref{fig:Li}.  The upper limit of A(Li) is obtained in another way.  Application of the method described in \\citet{tak05} using the averaged FWHM of weak Fe lines, and signal-to-noise ratio, to our spectra gives the upper limit of the equivalent width of this Li absorption EW(Li), EW(Li) $<$ 4.5 m\\AA\\ for both of our targets.  The upper limit of A(Li) is estimated to be A(Li) = 1.37 for KIC 9766237, and A(Li) = 1.43 for KIC 9944137, by calculations of the chemical abundances on the assumption of the local thermal equilibrium, atmospheric parameters obtained above, and that upper limit of EW(Li).  We here adopt the A(Li) upper limit of A(Li) = 1.0, that is, A(Li) $<$ 1.0 for our targets.  This value is the same as, or less than that of the Sun [A(Li)=0.92; \\cite{tak10}], and is relatively small in the range of the effective temperature around 5700 K \\citep{tak07}.  This fact suggests that these stars are not young (e.g. \\cite{dov14}), and both stars have an age greater than 3-4 Gy, unless an anomalous lithium depletion has occurred.  This is consistent with that both stars have a long rotation period, while their activity level seems to be higher than that of the Sun, based on the chromospheric data and on the lower limit for the relative spot areas of about 0.1 \\%, which is close to the largest sunspots groups observed (0.2 \\%).  \\begin{figure} \\begin{center} \\FigureFile(80mm,114mm){vlc-vsini_nogami-1.eps} \\end{center} \\caption{Diagram of the projected rotational velocity vs. rotational velocity estimated from the light curve.  The inclination angle is suggested to be fairly high, 90$^{+0}_{-25}$ deg for KIC 9766237, and 75$^{+15}_{-22}$ deg for KIC 9944137.  The horizontal error bar is estimated by taking into account the error of the period analysis ($\\sim$10 \\%), and the error of the radius estimation ($\\sim$10 \\%), and the vertical error bar is estimated on an assumption that the error is dominated by the error of the estimation of the macroturbulence ($\\sim$15 \\%; see \\cite{hir12}).} \\label{fig:vsini} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,114mm){Li6708-nogami-2.eps} \\end{center} \\caption{Li \\textsc{i} 6708 absorption line of KIC 9766237, KIC 9944137, and 18 Sco.  The solid lines represent synthetic spectra calculated for the stellar parameters listed in table \\ref{tab:param}, and different A(Li) values.  The Li \\textsc{i} 6708 line basically become shallower with a smaller Li abundance.  The A(Li) value of both targets is estimated to be A(Li) $<$ 1.0, although that of 18 Sco is estimated to be A(Li)=1.6($\\pm$0.2), in good agreement with the values in the literature.  The spectrum of KIC 9766237, and KIC 9944137 is shifted by +0.075, and +0.15, respectively, along the Y axis, for clarity.} \\label{fig:Li} \\end{figure}  We here briefly discuss whether superflares can be explained by the magnetic energy stored on the surface of KIC 9766237, and KIC 9944137. The average magnetic field of 1-20 G is deduced above, and the total energy of the superflare of our targets is 10$^{34}$ erg.  On the assumption that the radiated flare energy is a fraction [$f$; $f=0.1$ is assumed in \\citet{shi13a}] of the magnetic energy stored around the spot area, the flare energy ($E_{\\rm flare}$) is expressed as below (see \\cite{shi13a}): \\begin{equation} E_{\\rm flare} \\sim f \\frac{B^2}{8\\pi}L^{3} \\sim f \\frac{B^2}{8\\pi}(a\\pi R_*^2)^{3/2}. \\end{equation} In this equation, the length of the magnetic structure responsible for the flare, $L$, has been assumed to be equal to the size of the spotted region, i.e. $L=\\sqrt{a*\\pi*R_*^2}$, where $a$ is the area of the spot in units of the stellar disk area. This equation is transformed by replacing the energy with 10$^{34}$ erg, as below: \\begin{equation} a \\sim 0.01 \\Biggl( \\frac{f}{0.1} \\Biggr) ^{-\\frac{2}{3}} \\Biggl( \\frac{R_*}{\\RO} \\Biggr)^{-2} \\Biggl( \\frac{B}{10^3\\ {\\rm G}}  \\Biggr)^{-\\frac{4}{3}}. \\label{eq:spotarea} \\end{equation} If the strength of the magnetic field in the spot area is 20 kG, the fraction $a$ is deduced to be of the order of 0.001 with equation \\ref{eq:spotarea}.  In this case, the average magnetic field results in being $\\sim$20 G, which is in the range derived above.  The 20-kG magnetic field, however, should be too strong, compared to that of the Sun, and physically unrealistic, since the magnetic pressure of the spot area should be up to the same order of the gas pressure in the photosphere.  If the strength of the magnetic field in the spot area is 1 kG, the fraction $a$ is deduced to be of the order of 0.01.  In this case, the average magnetic field is of the order of 10 G, which is also in the range derived above.  The fraction a=0.01 is, however, an order of magnitude larger than that evaluated from the light curve above.  This implies the existence of a large polar spot which are always visible, and does not effectively contribute to the observed variability (see e.g. \\cite{hus07}).  More observations are needed for direct measurement of the spot group area, and the magnetic field strength of the spot group area.  In conclusion, although KIC 9766237, and KIC 9944137 are superflare stars having showed a superflare with an emitted energy of 10$^{34}$ erg, these stars turned out to be a G-type main sequence star with stellar parameters ($T_{\\rm eff}$, $\\log g$, and [Fe/H]) close to those of the Sun, by our high dispersion spectroscopy. Their slow rotation period of 21.8 days (KIC 9766237), and 25.3 days (KIC9944137), and very low lithium abundance of A(Li)$<$1.0 indicate that these stars are not young.  The projected rotational velocity derived here is consistent with the photometric rotation period, and these data indicate that the inclination angles are high (70-90 deg).  Line profiles of the photospheric absorption lines are in good agreement with those of 18 Sco, suggesting that both of the targets are single stars, though a low-mass companion can not be excluded without additional observations.  The average strength of the magnetic field on KIC 9766237, and KIC 9944137 is estimated to be 1-20 G, by the depth of Ca II 8542.  These arguments support the hypothesis that the Sun might cause superflares, though the upper limit of this value (20 G) is one order of magnitude larger than that of the Sun.  More detailed observations will be needed to clarify the magnetic field on the superflare stars.  \\bigskip  The authors sincerely thank the referee, Antonio Frasca for his very useful and constructive comments.  This study is based on observational data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.  We are grateful to Akito Tajitsu and other staffs of the Subaru Telescope for making large contributions in carrying out our observation.  Thanks are also to Youichi Takeda for his many useful advices on the analysis of our Subaru data, and for his opening the TGVIT, and SPTOOL programs into public.  We are indebted to Taichi Kato for his comments on S For.  Kepler was selected as the tenth Discovery mission.  Funding for this mission is provided by the NASA Science Mission Directorate. The Kepler data presented in this paper were obtained from the Multimission Archive at STScI. This work was supported by the Grant-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology of Japan (No. 25287039)."
    },
    "1402/1402.3258_arXiv.txt": {
        "abstract": "Models of rotating relativistic stars with a toroidal magnetic field have been computed for a sample of eight equations of state of cold dense matter. Non-rotating models admit important levels of magnetization and quadrupole distortion accompanied by a seemingly unlimited growth in size. Rotating models reach the mass-shedding limit at smaller angular velocities than in the non-magnetized case according to the larger circumferential equatorial radius induced by the magnetic field. Moreover, they can be classified as prolate--prolate, oblate--prolate, or oblate--oblate with respect to surface deformation and quadrupole distortion. Simple expressions for surface and quadrupole deformation are provided that are valid up to magnetar field strengths and rapid rotation. ",
        "introduction": "\\label{sec:intro}  Neutron stars with a strong toroidal magnetic field have attracted increasing interest as the magnetically induced distortion of their matter distribution may lead to the quasi-periodic emission of gravitational waves~\\cite{Cutler2002,Bonazzola1996}, for example, in the case of low-mass X-ray binaries (LMXBs). Moreover, strong magnetic fields are believed to power the electromagnetic activity of magnetars, which subsume both anomalous X-ray pulsars (AXPs) and soft-gamma repeaters (SGRs)~\\cite{Duncan1992,Thompson1996}. Models of relativistic stars with a toroidal magnetic field can be obtained within the standard formalism for stationary and axisymmetric relativistic stars~\\cite{Bonazzola1993}, since the electromagnetic stress--energy tensor then satisfies the same compatibility condition~\\cite{Oron2002} as the stress--energy tensor of an unmagnetized perfect fluid in purely rotational motion. Based on this finding, numerical models of relativistic stars with a toroidal magnetic field have emerged~\\cite{Kiuchi2008,Frieben2012} whereas the poloidal case was already studied a long time ago~\\cite{Bocquet1995}. ",
        "conclusions": ""
    },
    "1402/1402.3599_arXiv.txt": {
        "abstract": "Using deep narrow-band $H_2S1$ and $K_{\\rm s}$-band imaging data obtained with CFHT/WIRCam, we identify a sample of 56 \\Ha\\ emission-line galaxies (ELGs) at $z=2.24$ with the 5$\\sigma$ depths of $H_2S1=22.8$ and $K_{\\rm s}=24.8$ (AB) over 383 arcmin$^{2}$ area in the Extended Chandra Deep Field South. A detailed analysis is carried out with existing multi-wavelength data in this field. Three of the 56 \\Ha\\ ELGs are detected in \\textit{Chandra} 4\\,Ms X-ray observations and two of them are classified as active galactic nuclei. The rest-frame UV and optical morphologies revealed by \\textit{HST}/ACS and WFC3 deep images show that nearly half of the \\Ha\\ ELGs are either merging systems or have a close companion, indicating that the merging/interacting processes play a key role in regulating star formation at cosmic epoch $z=2-3$; About 14\\% are too faint to be resolved in the rest-frame UV morphology due to high dust extinction. We estimate dust extinction from spectral energy distributions. We find that dust extinction is generally correlated with \\Ha\\ luminosity and stellar mass. Our results suggest that \\Ha\\ ELGs are representative of star-forming galaxies. Applying extinction corrections for individual objects, we examine the intrinsic \\Ha\\ luminosity function (LF) at $z=2.24$, obtaining a best-fit Schechter function characterized by a faint-end slope of $\\alpha=-1.3$. This is shallower than the typical slope of $\\alpha=\\sim -1.6$ in previous works based on constant extinction correction. We demonstrate that this difference is mainly due to the different extinction corrections. The proper extinction correction is thus the key to recovering the intrinsic LF as the extinction globally increases with \\Ha\\ luminosity. Moreover, we find that our \\Ha\\ LF mirrors the stellar mass function of star-forming galaxies at the same cosmic epoch. This finding indeed reflects the tight correlation between star formation rate and stellar mass for the star-forming galaxies, i.e., the so-called main sequence. ",
        "introduction": "\\label{s:introduction}  The cosmic star formation rate (SFR) density peaks at $z\\sim 2-3$ and rapidly declines to the present day \\citep[][and references therein]{Hopkins06,Karim11,Sobral13}. Systematic studies of galaxy populations at the peak epoch hold the key to our understanding of galaxy formation and evolution, in particular for massive galaxies. The population of massive quiescent galaxies began to emerge \\citep[e.g.,][]{Brammer11} and the bulk of stars in local massive galaxies were formed at $z>\\sim 1.5$ \\citep{Renzini06}. A number of techniques have been developed to probe different galaxy populations in this crucial epoch, including Lyman break selection \\citep{Steidel99}, red color cut \\citep{Franx03}, $BzK$ selection \\citep{Daddi04}, submillimeter detection \\citep{Chapman05} and narrow-band imaging (\\citealt{Moorwood00}; see \\citealt{Shapley11} for a review). A complete picture is still missing due to the short of spectroscopic follow ups in the ``redshift desert'' $1.4<z<2.5$ for large samples selected in terms of physical quantities such as stellar mass or SFR.  The deep near-infrared (NIR) observations play a central role in probing high-$z$ galaxies in the rest-frame optical, which is more reflective of older stars (hence stellar mass). The NIR narrow-band imaging has turned out to be a modest way to identify emission-line galaxies (ELGs) in a narrow redshift range ($\\delta z/(1+z)=1\\%-2$\\%) over large sky coverage \\citep{Sobral13}. It provides redshifts with higher precision than the photometric redshift method based on broad-band spectral energy distributions (SEDs), and allows us to distinguish environments (e.g., groups or clusters) traced by the ELGs \\citep[e.g.,][]{Matsuda11,Hatch11,Geach12}. Moreover, the strength of an emission line, such as \\Ha\\ and \\OII $\\lambda$3727, can be measured by the flux excess in the narrow-band with respect to the corresponding broad-band. The optical and NIR emission lines from ionized gas surrounding massive young stars represent a nearly instantaneous measure of the SFR \\citep{Kennicutt12}. The \\Ha\\ is often used as an SFR indicator because it is  the strongest emission line in the optical/NIR, and less affected by dust obscuration compared to emission lines in the UV such as Ly$\\alpha$ in typical star-forming galaxies (SFGs).  The past decade has seen a number of NIR narrow-band surveys using ground-based telescopes \\citep[e.g.,][]{Geach08,Ly11} and NIR grism spectroscopic surveys with the \\textit{Hubble Space Telescope} \\citep[\\textit{HST}; e.g.,][]{Yan99,Atek10,vanDokkum11,Colbert13} in the study of \\Ha\\ ELGs at $z>$0.4, yielding \\Ha\\ luminosity functions (LFs) with a steep faint-end slope $\\alpha\\sim -1.6$ out to $z\\sim 2.5$ \\citep{Hayes10,Ly11,Sobral13}. The derived \\Ha\\ LFs suffer from large uncertainties mostly arising from poor extinction correction due to the lack of ancillary data. Detailed studies of the properties of \\Ha\\ ELGs will help to reduce the uncertainties. On the other hand, SFGs are found to follow a tight correlation between SFR and stellar mass \\citep[e.g.,][]{Noeske07,Daddi07}. This correlation, namely the main sequence, has been convincingly established from the local universe \\citep{Brinchmann04} to high-redshift universe \\citep{Guo13,Peng10,Karim11,Reddy12}. One would expect that the SFR function traces the stellar mass function (SMF) for SFGs. Since \\Ha\\ luminosity is a direct measure of SFR, the \\Ha\\ LF is expected to mirror the SMF. However, the SMF tends to have a shallow faint-end with a slope $\\alpha\\sim -1.3$ over a wide redshift range $0<z<3$ \\citep[e.g.,][]{Ilbert10,Brammer11}. The cause of this discrepancy in the faint-end slope between the two functions remains to be explored.  We carry out a deep 2.13\\,$\\micron$ narrow-band imaging survey to search for \\Ha\\ ELGs at $z$ = 2.24 in the Extended Chandra Deep Field South (ECDFS). This field contains the deepest observations from \\textit{Chandra}, \\textit{Galaxy Evolution Explorer}, \\textit{HST}, and \\textit{Spitzer}. With these data, we are able to securely identify a sample of \\Ha\\ ELGs, perform a detailed analysis and determine the intrinsic \\Ha\\ LF. We show that resolving dusty SFGs with intrinsically luminous \\Ha\\ from those with observed faint \\Ha\\ leads to a shallow faint-end slope for the \\Ha\\ LF. We describe observations and data reduction in Section~\\ref{s:observation}. Section~\\ref{s:Sample} gives the selection of \\Ha\\ ELGs and SED analysis to obtain photometric redshift, stellar mass and extinction. We present the properties of \\Ha\\ ELGs in Section~\\ref{s:property}. The derived \\Ha\\ LF at $z=2.24$ is given in Section~\\ref{s:Results}. We discuss and summarize our results in Section~\\ref{s:discussion}. Throughout the paper we adopt a cosmology with [$\\Omega_{\\Lambda}$, $\\Omega_M$, $h_{70}$] = [0.7, 0.3, 1.0]. Kroupa initial mass function \\citep[IMF; ][]{Kroupa01} and the AB magnitude system \\citep{Oke74} are used unless otherwise stated. ",
        "conclusions": "\\label{s:discussion}  We have described a search for \\Ha\\ ELGs using deep $H_2S1$ imaging combined with deep $K_\\mathrm{s}$ imaging data of ECDFS. In total 140 emission-line candidates are identified. The deep optical and NIR imaging data in the ECDFS are utilized to construct broad-band SEDs and derive photo-$z$, stellar mass and extinction. We identify 56 \\Ha\\ emitters with \\Ha\\ flux $>2.1\\times 10^{-17}$\\,erg\\,s$^{-1}$\\,cm$^{-1}$\\,Hz$^{-1}$. Three of the 56 are X-ray sources detected in the \\textit{Chandra} 4\\,Ms observation and two of them host an AGN in terms of their X-ray luminosities. The stacked X-ray luminosity of the rest is consistent with that of SFGs. In our sample, only $\\sim$ 4\\% are contaminated by AGN, although obscured AGNs individually undetected in the current X-ray observation may still present in some of the sample galaxies. However, the obscured AGNs are unlikely to be responsible for exciting \\Ha\\ line. As a result, our sample is clean for studying the properties of \\Ha\\ ELGs and determining the \\Ha\\ LF.  \\textit{HST}/ACS imaging reveals that nearly half of our \\Ha\\ ELGs are either mergers or close pairs. A similar merger fraction was suggested by \\citet{Conselice03} for SFGs at $z\\sim 2$, although large uncertainties exist in measuring merger fraction as a function of redshift \\citep[e.g., $\\sim$30\\% for \\Ha\\ ELGs at $z=2.23$ by][]{Stott13}. A significant fraction (14\\%) of the \\Ha\\ ELGs are too faint to be resolved in the rest-frame UV. These UV-faint ones are confirmed to be heavily-obscured SFGs by SED modeling. The remaining one third exhibit extended or compact morphologies. We note that the morphological properties of our \\Ha\\ ELGs are consistent with those of high-$z$ SFGs \\citep[e.g.,][]{Lotz06,Ravindranath06,Conselice11}. Our results reveal that merging/interacting processes play a key role in regulating star formation in typical SFGs at cosmic epoch $z\\sim 2-3$.  Our sample of 56 \\Ha\\ ELGs are representative of SFGs at $z\\sim 2-3$ in many aspects. SFGs are found to be more obscured at higher SFR \\citep[e.g.,][for $z\\sim 0$; \\citealt{Zheng07,Ly12,Dominguez12} for high-$z$]{Bothwell11}. We also find such a tendency among \\Ha\\ ELGs as shown in Figure~\\ref{f:atten_sfr.eps}. Our extinction is estimated from SEDs. It is important to note that the SED-derived extinction tends to be lower than the true case represented by the IR to UV luminosity ratio $L_\\mathrm{IR}/L_\\mathrm{UV}$, especially at the high-SFR end \\citep{Wuyts11a}. This is understandable because in starbursts, parts of newly formed stars are deeply buried in dusty molecular clouds and can only be traced by the IR emission \\citep[e.g.,][]{Swinbank04,Sedgwick13}. Adding this correction would enhance the correlation between extinction and SFR. The mass-SFR relation is well-established from the present day all the way to $z\\sim 3$ \\citep[e.g.,][]{Elbaz07,Daddi07}. The slope of the mass-SFR relation at $z\\sim 2$ is close to unity \\citep{Wuyts11b}. As shown in Figure~\\ref{f:main_sequence.eps}, \\Ha\\ ELGs follow the mass-SFR relation of SFGs, giving a slope of 0.81. We argue that the shallower slope of the \\Ha\\ ELGs is likely attributed to the underestimate of extinction at high end.  The correction for extinction on an individual basis is vital to determining the shape of the \\Ha\\ LF. With extinction correction for individual sample galaxies, we build the \\Ha\\ LF at $z=2.24$. The completeness as a function of \\Ha\\ luminosity is carefully estimated via Monte Carlo simulation and is applied to the observed LF. Our sample is not large and unable to give a good constraint on the bright-end of \\Ha\\ LF. A shallow faint-end slope of $\\alpha=-1.32\\pm 0.24$ is robust obtained, compared to the typical value of $\\sim -1.6$ given by previous studies in which correction for a constant extinction is often adopted \\citep{Geach08,Hayes10,Sobral13}. We have shown that a steep faint-end slope is inevitable when a fixed extinction is applied. \\citet{Garn10} computed the \\Ha\\ LFs at $z=0.84$ with correction for SFR-dependent extinction and constant ($A({\\rm H\\alpha})=1$\\,mag) extinction, respectively, finding that the former gives a shallower faint-end slope. It is obvious that the constant extinction correction is not a good approximation because it will overestimate the extinction in faint \\Ha\\ ELGs and severely underestimate the extinction for brightest ones. Resolving heavily-attenuated SFGs from the emitters with faint \\Ha\\ naturally leads to a decline at the faint end of an \\Ha\\ LF. The extinction is broadly correlated with SFR, stellar mass, and \\Ha\\ equivalent width \\citep[see also][]{Garn10,GarnBest10,Ly12,Dominguez12,Kashino13}. To address which parameter is more fundamental in governing extinction is beyond the scope of this work.  {\\it We argue that \\Ha\\ LF at $z=2.24$ mirrors the SMFs of SFGs at $z\\sim 2$}. We know that the SED-derived extinction tends to be underestimated. Therefore, our \\Ha\\ LF is likely underestimated too. Taking this into account, the agreement between our \\Ha\\ LF and these converted from SMFs is notably good. The evolution of SMFs or cosmic variance may also induce additional uncertainties. We emphasize that this finding is indeed a reflection of the main sequence of SFGs. The excellent agreement between the \\Ha\\ LF and the SMFs confirms that the main sequence has a slope close to unity. As a consequence, the \\Ha-selection and the color-selection target on the same population of SFGs and naturally link the two branches of studies together.  On the other hand, the faint-end slope of \\Ha\\ LF signifies the relative contribution between galaxies with weak and luminous \\Ha, a measurement that is important to the understanding of star formation behavior over cosmic time \\citep[e.g.,][]{Dutton10}. The faint-end of an LF is shaped by processes related to star formation, e.g., feedback from supernovae \\citep{Dekel86}, radiatively driven winds \\citep{Springel03}, and photons ionizing neutral gas \\citep{Kravtsov04}. The dependence of the faint-end slope of \\Ha\\ LF on extinction correction hints that UV LFs at high-$z$ might be in a similar situation. For instance, \\citet{Reddy09} presented the UV LF at $z\\sim 2-3$ with a faint-end slope of $-1.73$. Since UV light is closely coupled with \\Ha\\ emission, we suspect that the faint-end of the UV LF would become shallower if careful extinction correction is taken into account."
    },
    "1402/1402.1763_arXiv.txt": {
        "abstract": "In this work we propose a new matrix-free implementation of the Wiener sampler which is traditionally applied to high dimensional analysis when signal covariances are unknown. Specifically, the proposed method addresses the problem of jointly inferring a high dimensional signal and its corresponding covariance matrix from a set of observations. Our method implements a Gibbs sampling adaptation of the previously presented messenger approach, permitting to cast the complex multivariate inference problem into a sequence of uni-variate random processes. In this fashion, the traditional requirement of inverting high dimensional matrices is completely eliminated from the inference process, resulting in an efficient algorithm that is trivial to implement. Using cosmic large scale structure data as a showcase, we demonstrate the capabilities of our Gibbs sampling approach by performing a joint analysis of three dimensional density fields and corresponding power-spectra from Gaussian mock catalogues. These tests clearly demonstrate the ability of the algorithm to accurately provide measurements of the three dimensional density field and its power-spectrum and corresponding uncertainty quantification. Moreover, these tests reveal excellent numerical and statistical efficiency which will generally render the proposed algorithm a valuable addition to the toolbox of large scale Bayesian inference in cosmology and astrophysics. ",
        "introduction": "Ever increasing amounts and precision of modern cosmological and astrophysical data demands fast and robust methods to address corresponding large scale inference problems of analysing these observations and extracting new knowledge on our Universe. Particularly, the Wiener filter has become a standard tool for the analysis of large data sets, often involving many millions of parameters, with widespread applications in cosmology and astrophysics. Even though relying on a linear data model and Gaussian statistics, the Wiener filter approach is still a standard and well valued method when requiring a robust approach for the inference of high dimensional signals as occurring in the analysis of large scale structure (LSS) or cosmic microwave background (CMB) data. For this reason, Wiener filtering has been frequently applied to a variety of large scale structure analysis problems, specifically the inference of the three dimensional density field from galaxy observations \\citep[see e.g.][]{1994ApJ...423L..93L,WienerFSL,1995MNRAS.272..885F,1993ApJ...415L...5G,1995ApJ...449..446Z,1995MNRAS.272..885F,HOFFMAN1994,1995MNRAS.277..933S,1999ApJ...520..413Z,2002MNRAS.331..901Z,2004MNRAS.352..939E,KITAURA2008,2006MNRAS.373...45E,KITAURA2009,JASCHE2010PSPEC,JASCHE_SPEC2013}.  In the field of CMB analysis, the Wiener filter is frequently employed as a map making algorithm or for the joint inference of temperature fluctuations and corresponding power-spectra \\citep[see e.g.][]{2004ApJS..155..227E,JEWELL2004,ODWYER2004,WANDELT2004,2004ApJS..155..227E,SMITH2007,LARSON2007,ERIKSEN2007,ELSNER2013}. Similar approaches have also been used for the optimal reconstruction of images from radio Interferometry \\citet{2006ApJS..162..401S,SUTTER2013} or to generate improved maps of the galactic Faraday emission \\citet{2012A&A...542A..93O}. Furthermore, the Wiener filter also constitutes an integral part of the recently presented general purpose statistical analysis framework NIFTY (Numerical Information Field Theory) \\citep[][]{SELIG2013}.  Traditionally, numerical implementations of the Wiener filter rely on Krylov space methods, such as conjugate gradients, to invert matrices and solve high dimensional systems of linear equations \\citep[see e.g.][and references therein]{KITAURA2008}. However, implementation and testing of these numerically demanding methods also constitutes a certain hurdle and generally requires some investment into code development before such methods can be used for specific scientific applications.  Recently a particular elegant and simple way of implementing the Wiener filter via a messenger method has been proposed, which remedies the hurdle of implementing numerical matrix inversion techniques for very high dimensional systems \\citep[][]{ELSNER2013}. In particular, \\citet{ELSNER2013} proposes to introduce a messenger field to mediate between different preferred orthogonal bases, in which signal and noise covariance matrices can be expressed conveniently. In this approach information from the data is transmitted to the signal via a messenger field, that can generally be transformed efficiently from one bases representation to another. In this fashion the algorithm avoids the requirement to apply the inverse Wiener covariance matrix to data.  In this work, we will pick up these ideas and propose an efficient and easy to implement Gibbs sampling approach to address Bayesian large scale inference problems in cosmology or astrophysics.  Our primary intention is to propose a simple, nevertheless powerful, algorithm for the joint inference of a signal and corresponding covariance matrix from observations that can be implemented and operated by everyone, even inexperienced users. Specifically, in this work, we will exemplify the performance of this algorithm in case of a LSS analysis aiming at the joint inference of the three dimensional density field and cosmological power-spectrum from galaxy surveys.  The introduction of a messenger field yields an augmented Wiener posterior distribution, whose structure lends itself to a ideal multiple block sampling approach. In a first step the messenger field is realised from a normal distribution conditional on the observation in real space, followed by a second step of sampling the signal conditional on the previously sampled messenger field in Fourier-space. Here, the augmented Wiener posterior distribution is chosen such, that sampling these two random fields can be trivially achieved by generating a sequence of uni-variate normal random variates in their respective basis representations. Iterating these processes will then provide samples from the Wiener posterior without the need of performing matrix inversions or any other multi-parameter operation, except for basis transformations. Further we will complement this algorithm, by a power-spectrum sampling method to jointly infer the signal and its covariance matrix. Likewise, as described in the following, this algorithm will also only require the ability to generate uni-variate normal and inverse gamma variates.  Consequently, we arrive at an efficient algorithm, which, at every stage, reduces the full joint problem to a sequence of independent uni-variate sub-problems. The advantage of this algorithm lies in its ease of implementation, thus greatly reducing the required investment in code development for large scale Bayesian inference projects. In the following we will describe the implementation of this algorithm and exemplify it in case of a mock LSS analysis. The paper is structured as follows. In section \\ref{augmented_posterior} we will describe the messenger approach and the resulting augmented Wiener posterior distribution. Section \\ref{implementation}, describes the numerical implementation of the proposed method. To estimate the performance of the algorithm in a realistic scenario, as an example, we will apply it to an artificial galaxy survey, which will be described in section \\ref{mock_Data}. In the following section \\ref{Results}, we will discuss the results of these tests and conclude the paper with a summary and conclusion in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} Modern cosmology has an ever increasing demand for fast and accurate statistical inference methods to counter present and upcoming avalanches of cosmological and astrophysical data. As pointed out in the introduction, inference of signals from observations subject to noise is a ill-posed problem requiring sophisticated statistical methods to quantify corresponding statistical uncertainties. Specifically large scale Bayesian inference, such as the joint inference of three dimensional matter density fields and corresponding cosmological power-spectra, from observations relies on complex and numerically expensive MCMC methods, often involving implementations of Krylov space methods or gradient based Hybrid Monte Carlo approaches \\citep[see e.g.][]{KITAURA2008,JASCHE_SPEC2013}. Not only are these methods numerically expensive but are also hard to test and debug especially in large scale applications \\citep[][]{COOK2006}. Besides these issues, the requirement to implement Krylov space or HMC methods constitutes a significant hurdle for rapid prototyping and development of large scale Bayesian inference methods in cosmology and astrophysics. To address these issues, in this work we present a new, efficient and trivial to implement Gibbs sampling approach for the joint inference of cosmological density fields and power-spectra for linear data models. This approach picks up basic ideas of the recently proposed messenger method for Wiener filtering, described by \\citet{ELSNER2013}, and does not require any matrix inversions to explore high dimensional parameter spaces. As described in section \\ref{augmented_posterior}, introducing a messenger field to mediate between different preferred bases, in which signal and noise covariance matrices can be expressed conveniently, yields an augmented Wiener posterior distribution which can be efficiently explored via a multiple block Gibbs sampling approach \\citep{ELSNER2013}. In particular, the proposed method turns the cumbersome approach of inverting multi-million dimensional matrices  into the task of sequentially drawing random numbers from only uni-variate normal distributions. While trivial to implement, iteration of this process yields full multi-variate random fields drawn from the desired target Wiener posterior distribution, hence correctly addressing the large scale inference problem. To address also problems in which the covariance matrix of the signal to infer is unknown, we add a power-spectrum block sampler to jointly infer the signal and it's power spectrum.  As described in section \\ref{implementation}, this power-spectrum sampler permits to efficiently explore the high as well as the low signal to noise regime. While low signal to noise regimes are dominated by stochastic noise, high signal to noise regions are dominated solely by the much smaller sample variance. Using this sample variance to determine step sizes in Markov transitions will result in correct sampling of the target posterior, but may also yield unfeasibly long Markov chains as parameters in the low signal-to-noise regime remain correlated over many steps. To counter this problem, in \\ref{implementation} we introduced a coordinate transform which permits to perform larger steps in low-signal-to noise regimes via a Metropolis Hastings transition step. The combination of both approaches yields a covariance matrix sampler that is efficient at all regimes, while only requiring the ability to generate uni-variate random numbers.  In this fashion the task of jointly sampling a signal and its covariance matrix can be addressed purely by a sequence of uni-variate sampling processes.  In section \\ref{Results} we exemplify the performance of our method in a cosmological setting by applying it to a artificial galaxy mock catalogue, described previously in section \\ref{mock_Data}, aiming at the joint inference of the three dimensional density distribution and its cosmological power-spectrum from observations. This artificial data set emulates dominant features of the Sloan Digital Sky Survey data release 7, in particular survey geometry, selection effects and noise, and thus constitutes a realistic test scenario.  A particular important aspect, when dealing with Markov Chain Monte Carlo algorithms, is the determination of their statistical efficiency. As any Markov Chain Monte Carlo method generates a sequence of correlated samples the amount of actually produced independent samples is limited by the total length of the chain. In section \\ref{stat_eff} we therefore analysed the intra-chain correlation length between subsequently generated samples of the cosmological power-spectrum. These test demonstrates formidable statistical efficiency for the proposed method over the entire range of Fourier-modes present in the analysis. Specifically these tests indicate that the proposed Markov algorithm generates independent samples at every 50{\\it{th}} iteration of the Markov chain, where we chose one cycle to consists in ten Markov transitions.  Section \\ref{inf_spec_dens} discusses the results obtained by the proposed Markov method. In particular, the method provides estimates for the ensemble mean cosmological power-spectrum and corresponding uncertainty quantification. The inferred power-spectrum recovers the underlying true signal and shows no sign of bias throughout the entire range of Fourier-modes under consideration. These result are also consistent with previous results \\citep[see ][]{JASCHE2010PSPEC,JASCHE_SPEC2013}.  Furthermore, in our example case, the method also provides inferred three dimensional maps of the cosmic matter distribution. In Figure \\ref{fig:mean_var_dens} we demonstrate ensemble mean estimates of the density field and ensemble covariance maps quantifying the corresponding uncertainty. The proposed method therefore not only provides single estimates of signals, but also provides means to quantify and propagate statistical uncertainties for any finally inferred quantity, as is required for modern precision cosmology.  The ease of implementation, numerical and statistical efficiency renders this method an ideal tool for large scale Bayesian applications involving million dimensional problems and linear data models. A particularly important feature of the method is, that it only requires the ability to sample from uni-variate distributions and thus can be trivially implemented and tested by even inexperienced users or can be used for rapid prototyping and development of more complex inference frameworks.  In summary, we propose a statistical and numerically efficient Gibbs sampling approach for the inference of a unknown signal and its covariance matrix from observations subject to statistical uncertainties and systematics. Particularly due to the ease of implementation we anticipate this method to greatly add to the propagation of high precision large scale data analysis methods in cosmology and astrophysics, eventually leading to a more complete understanding of our Universe."
    },
    "1402/1402.1280_arXiv.txt": {
        "abstract": "We present results of a \\Chandra High Energy Transmission Grating (HETG) observation of the carbon-rich Wolf-Rayet (WR) star \\WRE. These are the first high-resolution spectra of this object in X-rays. Blue-shifted centroids of the spectral lines of $\\sim -360$\\kms and line widths of $1000 - 1500$\\kms (FWHM) were deduced from the analysis of the line profiles of strong emission lines. The forbidden line of Si XIII is strong and not suppressed, indicating that the rarified 10-30 MK plasma forms far from strong sources of far-UV emission, most likely in a wind collision zone. Global spectral modeling showed that the X-ray spectrum of \\WR suffered higher absorption in the October 2012  \\Chandra observation compared to a previous  January 2008  \\XMM observation. The emission measure of the hot plasma in \\WR decreased by a factor $\\sim 3$ over the same period of time. The most likely physical picture that emerges from the analysis of the available X-ray data is that of colliding stellar winds in a wide binary system with an elliptical orbit. We propose that the unseen secondary star in the system is another WR star or perhaps a luminous blue variable. ",
        "introduction": "\\WR is a carbon-rich (WC) Wolf-Rayet  star with a WC8 spectral classification \\citep{vdh_01} that was discovered in a near-infrared survey by \\citet{danks_83}. This WC star is located inside the G305 star-forming region in the Scutum Crux arm of the Galaxy. Its proximity (within $2'$) to the two compact infrared clusters Danks 1 and 2 suggests that \\WR likely originates from one or the other \\citep{danks_84}. The optical extinction toward \\WR is very high, A$_V = 9.2$~mag \\citep{danks_83}, and only a small part of it is due to circumstellar material \\citep{baume_09}. The distance to \\WR is not yet well-constrained and various studies provide a range of $1.21 - 4$~kpc  (e.g., \\citealt{vdh_01}; \\citealt{baume_09}; \\citealt{danks_83}).  The evolution of its infrared emission suggests that \\WR is a long-period  episodic dust-maker \\citep{williams_95}. This was confirmed by a recent study that revealed  recurrent dust formation on a time scale of more than 32 years which also indicates that \\WR is very likely a wide colliding-wind binary \\citep{williams_12}. However,  the interpretation  of the nature of \\WR may not be straightforward. Recently, \\citet{hindson_12} reported the detection of a {\\it thermal} radio source (spectral index $\\alpha = 0.6$, F$_{\\nu} \\propto \\nu^{\\alpha}$) associated with \\WRE, while as a rule  wide colliding-wind binaries are {\\it non-thermal} radio sources \\citep{do_00}.   Nevertheless, the binary nature of \\WR is further supported by its X-ray characteristics. Analysis of the \\XMM spectra of \\WR showed that its X-ray emission is of thermal origin and this is the most X-ray luminous WR star in the Galaxy detected so far, after the black-hole candidate Cyg X-3 \\citep{zhgsk_11}. The latter is valid provided \\WR is associated with the open clusters Danks 1 and 2 and is thus located at the distance of $\\sim 4$~kpc \\citep{danks_83}.  It is important to recall that all the previous pointed X-ray observations of a small sample of presumed single WC stars have yielded only non-detections, demonstrating that they are X-ray faint or X-ray quiet (\\citealt{os_03}; \\citealt{sk_06}). Therefore, the high X-ray luminosity of the WC star \\WR is a clear sign that it is not a single star and it is very likely that its enhanced X-ray emission originates from the interaction region of the winds of two massive binary components (\\citealt{pril_76}; \\citealt{cherep_76}).  In this paper, we report results from the first grating observation of the dusty WR star \\WRE. In Section \\ref{sec:data}, we briefly review the \\Chandra HETG observation. In Section \\ref{sec:overall}, we present an overview of the X-ray spectra. In Section \\ref{sec:lines}, we analyze the profiles and ratios of strong X-ray emission lines. In Section \\ref{sec:fits}, we present the results from the global spectral fits. We discuss the results from our analysis in Section \\ref{sec:discussion} and list our conclusions in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In this work, we presented an analysis of the first grating X-ray spectra of the dusty WR star \\WR obtained with the \\Chandra HETG. The basic results and conclusions are as follows.  \\begin{enumerate} \\item From analysis of the line profiles of strong emission lines, a typical line width (FWHM) of $1000 - 1500$\\kms was deduced and blue-shifted line centroids of $\\sim -360$\\kms. A strong (not suppressed) forbidden line in the He-like triplet of Si XIII was detected which indicates that this line forms a rarefied hot plasma far from strong sources of UV emission.  \\item Global spectral modeling showed that the X-ray spectrum of \\WR suffered higher absorption (likely of {\\it wind} origin) in October 2012 (the \\Chandra observation) compared to January 2008 (the \\XMM observation). The emission measure of the hot plasma in \\WR decreased by a factor $\\sim 3$ over the same period of time.  \\item No X-ray variability on a timescale of less than 100 ksec was detected. This result is similar to what was found from the analysis of the previous X-ray observation of \\WR with \\XMM \\citep{zhgsk_11}.  \\item The most likely physical picture that emerges from the analysis of the available X-ray data (\\Chandra and \\XMM) is the following. The high X-ray luminosity of the carbon-rich WR star \\WR is due to colliding stellar winds in a wide binary system with elliptical orbit. The observed changes of the characteristics (emission measure, absorption) of the X-ray emission of \\WR and the blue-shift of the line centroids of the strong lines can find their place in the CSW picture, provided the secondary star in the binary system (the companion star of the primary WC object) is not an O star but instead a WR star or a LBV star.  \\item More X-ray observations with high spectral resolution are needed to describe in detail the X-ray characteristics of \\WRE. Due to the high X-ray absorption of this object (most of the X-ray emission from \\WR emerges at energies above 2 keV), \\Chandra HETG data are essential for us to obtain a deeper understanding of the physical picture of CSWs in the fascinating dusty WR star \\WRE.   \\end{enumerate}"
    },
    "1402/1402.1249_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{introduction}  Solar flares represent a process of rapid transformation of the magnetic energy of active regions into the kinetic energy of charged particles, plasma flows and heating of the solar atmosphere and corona.  The primary energy release during the flares is believed to occur in the corona as a result of magnetic reconnection.  It is generally believed that most of the energy released by the reconnection goes directly and indirectly (via plasma waves) to acceleration of electrons and protons which are injected into flaring loops (Fig.~\\ref{fig1}$a$). Most of the observed radiation  is produced either directly by these particles or indirectly through energization of the background plasma.  \\begin{figure} \\begin{center} \\includegraphics[width =\\textwidth]{kosovichev_fig1.eps} \\end{center} \\caption{a) Illustration of a sunquake mechanism in the `thick-target' model of solar flares: high-energy particles injected into a magnetic flux tube heat an upper chromospheric layer, creating a high-pressure region which produces `chromospheric evaporation' and a shock traveling to the photosphere and causing a sunquake impact. b) Illustration of ray paths of acoustic waves excited by an impulsive source on the solar surface. The waves travel into the deep interior and are reflected back to the surface, producing sunquake ripples. \\label{fig1}} \\end{figure}  It was suggested long ago \\citep{Wolff1972} that flares may cause acoustic waves traveling through the solar interior, similar to the seismic waves in the Earth. Because the sound speed increases with depth, the waves are reflected in the deep layers of the Sun (Fig.~\\ref{fig1}$b$), and then appear on the surface, forming expanding rings, similarly to ripples on the water surface. Theoretical modeling \\citep{Kosovichev1995} predicted that the speed of the expanding seismic ripples increases with distance because the distant waves propagate into the deeper interior where the sound speed is higher.  First observations of the seismic waves caused by the X2.6 flare of 9 July, 1996,  proved these predictions \\citep[Fig.~\\ref{fig2};][]{Kosovichev1998}. These observations also showed that the source of the seismic response was a strong shock-like compression wave propagating downwards in the photosphere. This led to a suggestion that the seismic response can be explained in terms of so-called the hydrodynamic ``thick-target'' model (Fig.~\\ref{fig1}$a$). In this model  \\citep[e.g.][]{Kostiuk1975, Livshits1981, Fisher1985, Kosovichev1986}, a beam of high-energy particles, accelerated in the corona, heats the solar chromosphere, resulting in  a strong compression of the lower chromosphere. This  compression produces a chromospheric eruption (`evaporation') and a downward-propagating shock wave which hits the solar surface and causes a seismic response. This shock is observed in the SOHO/MDI Dopplergrams as a localized large-amplitude velocity impulse of about 1~km/s, immediately after the hard X-ray impulse produced by high-energy electrons (Fig.~\\ref{fig3}$a$). This velocity impulse represents the initial hydrodynamic impact exciting the seismic waves. In addition,  it was found that the seismic waves had a significant anizotropic quadrupole component. However, the seismic wave front was almost circular (Fig.~\\ref{fig2}).   \\begin{figure} \\begin{center} \\includegraphics[width =0.9\\textwidth]{kosovichev_fig2.eps} \\end{center} \\caption{\\textit{A sequence of enhanced Doppler-velocity images of the first sunquake observed on July 9, 1996, with the Michelson Doppler Imager on SOHO spacecraft. The sunquake is observed on the solar surface as circular expanding wave originated from an impact of a X2.6 solar flare. Dark color shows upflows, and white color shows downflows. The noisy background is caused by solar granulation. Bright permanent features are flows associated with magnetic field of the active region. The sunquake wave signal in these images is enhanced by a factor of 4.} \\label{fig2}} \\end{figure}   Figure~\\ref{fig3}$b$ shows a time-distance diagram for the spherically symmetrical component of the sunquake wave. The diagram is obtained by azimuthally averaging the Doppler velocity signal centered at the initial impact location. In this diagram, the initial impact started at about 9:10~UT, and the seismic wave is displayed as a narrow ridge with a characteristic slope, which follows the theoretical time-distance relations for acoustic waves traveling in the unperturbed conditions of the quit Sun. In this case, the theoretical relation is calculated for a standard solar model using the acoustic ray-path approximation \\citep[e.g.][]{Kosovichev2011b}. While the appearance of sunquakes is similar to water ripples their physical properties are quite different. In particular, as follows from Fig.~\\ref{fig3}$b$ the surface speed of sunquake waves rapidly increases with the travel distance. Immediately after the initiation the wave speed is about 10 km/s, which is close to the surface speed of sound, and then it rapidly increases to $\\sim 100$~km/s. The apparent surface speed increases because the acoustic waves travel through the solar  interior where the sound speed is much higher than on the surface, and more distant waves travel through deeper interior (see Fig.~\\ref{fig1}$b$).  \\begin{figure} \\begin{center} \\includegraphics[width =0.85\\textwidth]{kosovichev_fig3.eps} \\end{center} \\caption{a) The hard X-ray flux of the July 9, 1996, flare (solid) and the Doppler velocity signal of the localized flare impact which generated a sunquake; b) a time-distance diagram of the sunquake wave, obtained by azimuthally averaging and stacking the Doppler-velocity images obtained with 1-min cadence. The strong variations of velocity at about 9:10 UT represent the initial impact. The white dashed curve shows a theoretical time-distance relation of acoustic waves, calculated in a ray-path approximation. \\label{fig3}} \\end{figure}   Subsequent observations of solar flares made by the Michelson Doppler Imager (MDI) instrument on the NASA-ESA mission SOHO \\citep{Scherrer1995} did not show noticeable sunquake signals even for X-class flares until the big flares of October 2003. This search was carried out by calculating an``egression'' power for high-frequency acoustic waves during the flares \\citep{Donea1999}. Thus, it was suggested that sunquakes are a rather rare phenomenon on the Sun, which occur only under some special conditions. Surprisingly, seven years after the first event several flares did show strong ``egression'' signals indicating new potential sunquakes \\citep{Donea2005}. Of course, detection of sunquakes depends on their amplitude relative to the background solar noise. Presumably, all flares generate some seismic response, but if the amplitude is not high enough the waves may be lost in the background noise.  The helioseismology observations with MDI discovered several powerful events in 2003-2005 during the declining phase of Solar Cycle 22  \\citep{Kosovichev2006b,Donea2011}. Perhaps, the most powerful event of the cycle was on January 15, 2005 (Sec.~\\ref{thick-target}). The observations revealed that in most events the helioseismic response is anisotropic, and that the wavefront shape may significantly deviate from circular. These effect are mostly associated with the sunquake source extension along flare ribbons and also with the source motion which can reach supersonic speeds \\citep{Kosovichev2006a}. The source motion is often related to expanding flare ribbons, and probably reflects properties of magnetic reconnection and energy release in the solar atmosphere \\citep{Sturrock1966a}. Recent observations with the Helioseismic and Magnetic Imager instrument on board Solar Dynamics Observatory \\citep{Scherrer2012} show that sunquakes are much more common events than it was thought before, and that they can be frequently observed for moderate-class flares \\citep{Kosovichev2012a}.  In this article, I present a brief overview of the basic theory and observational properties of sunquakes, discuss some recent results. ",
        "conclusions": ""
    },
    "1402/1402.6719_arXiv.txt": {
        "abstract": "Using the Westerbork Synthesis Radio Telescope (WSRT), we obtained high-time-resolution measurements of the full (linear and circular) polarization of the Crab pulsar. Taken at a resolution of 1/8192 of the 34-ms pulse period (i.e., $4.1~\\mu{\\rm s}$), the 1.38-GHz linear-polarization measurements are in general agreement with previous lower-time-resolution 1.4-GHz measurements of linear polarization in the main pulse (MP), in the interpulse (IP), and in the low-frequency component (LFC). We find the MP and IP to be linearly polarized at about $24\\%$ and $21\\%$, with no discernible difference in polarization position angle. However, and contrary to theoretical expectations and measurements in the visible, we find no evidence for significant variation (sweep) in polarization position angle over the MP, the IP, or the LFC. Although, the main pulse exhibits a small but statistically significant quadratic variation in the degree of linear polarization. We discuss the  implications which appear to be in contradiction to theoretical expectations. In addition, we detect weak circular polarization in the main pulse and interpulse, and strong ($\\approx 20\\%$) circular polarization in the low-frequency component, which also exhibits very strong ($\\approx 98\\%$) linear polarization at a position angle about $40\\degree$ from that of the MP or IP. The pulse-mean polarization properties are consistent with the LFC being a low-altitude component and the MP and IP being high-altitude caustic components. Nevertheless, current models for the MP and IP emission do not readily account for the observed absence of pronounced polarization changes across the pulse. Finally, we measure IP and LFC pulse phases relative to the MP that are consistent with recent measurements, which have shown that the phases of these pulse components are evolving with time. ",
        "introduction": "The Australia Telescope National Facility Pulsar Catalog \\citep{Manc05} lists over 2300 radio pulsars. Several radio studies \\citep[e.g.,][]{Goul98,Kara06,Welt08} have measured the polarization for many of these pulsars. Radio pulsars typically show moderate-to-strong linear polarization ($p_{L}$), being stronger for those of higher spin-down energy-loss rate \\citep[Figure~8]{Welt08}. The linear polarization sometimes exhibits a characteristic swing or sweep of the position angle in an S-like shape near the pulse center, which is routinely interpreted in terms of the rotating vector model \\citep[RVM,][]{Radh69}. For this model the point of emission is assumed to be in the polar cap region of the pulsar where a dipolar magnetic-field line points with a small angle (beamwidth) towards the observer. The two free parameters of this simple model are the angle between the axes of rotation and the orientation of the magnetic dipole, and the view angle between the line of sight and the rotation axis. The variation of the radio position angle from some pulsars \\citep[e.g.,][and references therein]{Lyne06} can be described by this model.  The Crab pulsar, the compact remnant of SN1054, and its pulsar wind nebula (PWN) are amongst the most intensively studied objects in the sky. The pulsar is one of the youngest and most energetic and its pulsed emission has been detected from 10 MHz \\citep{Brid70} up to 400~GeV by VERITAS \\citep{Aliu11} and by MAGIC \\citep{Alek12}. The PWN is detected at energies up to 100~TeV \\citep{Ahar04,Ahar06,Alle07,Abdo12}. Both the pulsar and nebula are predominantly sources of non-thermal radiation (synchrotron, curvature, and Compton processes), indicated not only by the broadband spectral continua but also by strong polarization in many wavelength bands \\citep{Lyne06,Bueh13}.  In the visible band, spatially-resolved polarimetry of the nebula, which began over a half century ago \\citep{Oort56,Wolt57}, continues \\citep[e.g.,][and references therein]{Mora13b}. Owing to its brightness, phase-resolved optical polarimetry of the pulsar has also been possible \\citep{Jone81,Smit88,Slow09}. However, phase-resolved X- and $\\gamma$-ray polarimetry measurements of the Crab pulsar require space-based instruments, which have had limited sensitivity. OSO-8 observations \\citep{Silv78} of the Crab established only upper limits to the X-ray (2.6-keV and 5.2-keV) polarization of the pulsed emission. INTEGRAL IBIS observations \\citep{Foro08,Mora13a} also detect no significant pulsed $\\gamma$-ray (200--800-keV) polarization, although the off-pulse emission appears highly linearly polarized and is possibly associated with structures close to the pulsar rather than with the pulsar itself.  The Crab pulsar's light curve exhibits different features at different wavelengths, but it is currently the only pulsar for which the principal features persist over all wavelengths, from radio to $\\gamma$-ray. There are two principal components---the main pulse (MP) and the interpulse (IP). This double-peak structure remains more-or-less phase-aligned over all spectral bands \\citep{Moff96,Kuip01}. One of several additional features in the radio band is the low-frequency component (LFC, e.g., \\citeauthor{Moff96} \\citeyear{Moff96,Moff99}), having very low amplitude and occurring $\\approx 0.10$ fractional pulse phase ($36\\degree$) before the MP. This component is most prominent around 1.4~GHz, in contrast with the ``precursor'' component \\citep{Moff96}, which precedes the MP by $\\approx 0.05$ fractional pulse phase ($19\\degree$) at 0.327 and 0.610~GHz \\citep[Table 2 of][]{Backer00}.  The MP and IP appear at roughly the same pulse phase from radio to $\\gamma$-ray wavelengths, suggesting that their emission originates from a similar location in the magnetosphere at all wavebands. Modeling of $\\gamma$-ray light curves from the many pulsars observed by the {\\sl Fermi Gamma-ray Space Telescope} \\citep{Abdo13} strongly indicates that the high-energy emission originates in the outer magnetosphere, at altitudes comparable to the light-cylinder radius \\citep{Roma10, Pier13, BaiSpit10}. Outer magnetosphere emission models, such as the outer-gap \\citep{Roma95}, slot-gap \\citep{Musl04}, and current-sheet \\citep{Petr05} had been proposed and studied prior to the {\\sl Fermi} observations, but their emission geometry seems to account for the characteristics and variety of observed $\\gamma$-ray light curves. In addition, {\\sl Fermi} has discovered a number of $\\gamma$-ray millisecond pulsars whose radio peaks are nearly aligned with their $\\gamma$-ray peaks \\citep[e.g.,][]{Espi13}, like the Crab. Modeling both $\\gamma$-ray and radio light curves of these pulsars with the same outer magnetosphere emission models used to model young pulsars has suggested that their radio emission may originate from very high altitudes \\citep{Vent12}. Thus, in this paper we compare the phase-resolved radio polarization observations (\\S\\ref{s:obs}) that we have analyzed (\\S\\ref{s:res}) with such models (\\S\\ref{s:imp}).  \\citet{Manc71} measured the linear polarization of the Crab pulsar's MP and precursor components at two radio frequencies. The 0.410-GHz measurements found the MP to be $20\\%$ linearly polarized at position angle $140\\degree$ and the precursor to be $80\\%$ linearly polarized at position angle $140\\degree$. The 1.664-GHz measurements found the MP to be $24\\%$ linearly polarized at position angle $60\\degree$ and the precursor to be completely absent. As these measurements had rather large uncertainties and were obtained with a time resolution $1/256$ of the pulse period, they were quite limited for detecting variation of the linear polarization degree or position angle within a feature. However, \\citeauthor{Manc71} noted a suggestion of rotation of the 1.664-GHz polarization position angle by about $30\\degree$ through the MP.  More recently, \\citet{Moff99} examined the pulse-profile morphology and polarization properties at three radio frequencies---1.424~GHz, 4.885~GHz, and 8.435~GHz---with a time resolution of $256~\\mu{\\rm s}$ (about $1/130$ of the pulse period). The 1.424-GHz measurements found the MP to be $25\\%$ linearly polarized at position angle $120\\degree$; the IP, $20\\%$ at position angle $120\\degree$; and the LFC, $45\\%$ at position angle $155\\degree$. \\citeauthor{Moff99} note that the polarization position angle ``changes across the full period, although not significantly within components''.  Here we first report our observations (\\S\\ref{s:obs}), using the Westerbork Synthesis Radio Telescope (WSRT) in the Netherlands, of the full (linear and circular) 1.38-GHz polarization of the Crab pulsar, at high time resolution. We then describe the polarimetry analysis and results (\\S\\ref{s:res} and Appendix~\\ref{s:stat}) for the three pulse components, with a primary objective of determining the sweep of the position angle across each. Next we discuss the implications (\\S\\ref{s:imp}) of our measurements and analysis upon theoretical models for the pulsar emission. Finally, we summarize our conclusions (\\S\\ref{s:sum}). ",
        "conclusions": "\\label{s:sum}  Our 1.38-GHz observations of the Crab pulsar measured significant linear and circular polarization in the three most prominent pulse components---the main pulse (MP), inter pulse (IP), and low-frequency component (LFC). These results are mostly in agreement with previous measurements of linear polarization at similar radio frequencies \\citep[cf.][]{Moff99}. The MP and IP are moderately linearly polarized ($p_{L}\\approx 23\\%$ and $21\\%$, respectively) at the same position angle ($\\psi_{\\rm IP}-\\psi_{\\rm MP}\\approx 0$); they are weakly circularly polarized ($p_{C}\\approx -1.3\\%$ and $-3.2\\%$, respectively). In contrast, the LFC is very strongly linearly polarized ($p_{L}\\approx 98\\%$), at a position angle $+40\\degree$ from that of the MP or IP, and moderately circularly polarized ($p_{C}\\approx 20\\%$).  The fine time resolution (Period/8192 = 4.1 $\\mu$s) and good sensitivity of the measurements at the Westerbork Synthesis Radio Telescope (WSRT) enabled a meaningful search for changes in linear-polarization degree $p_{L}$, in position angle $\\psi$, and in circular-polarization degree $p_{C}$ across each of the three pulse components. Neither the MP, IP, nor LFC exhibits a statistically significant change in the polarization position angle or circular polarization across the pulse. For the MP, the linear term (``sweep'') is well constrained: $\\psi'_{0\\,\\rm MP}=(-0.16\\pm 0.20)\\degree{\\rm PA}/\\degree$. Likewise, neither the IP nor LFC displays a statistically significant change in the polarization degree. However, the MP does show a small but statistically significant quadratic variation in linear-polarization degree---$p''_{L0\\,\\rm MP}=(0.88\\pm 0.22)\\%/\\degree/\\degree$ about its central value---$p_{L0\\,\\rm MP}=(23.0\\pm 0.3)\\%$---for a pulse-average linear polarization $\\overline{p}_{L\\,\\rm MP}=(23.7\\pm 0.3)\\%$.  Our analysis of the radio Stokes data shows no strong sweep of the linear-polarization position angle. This lack of strong position-angle swings contrasts with the rapid swings observed in the visible band. Current models for pulsar emission geometries do not readily account for the absence of substantial variations in both polarization degree and position angle across a pulse component (\\S~\\ref{s:imp}). Thus, alternative models---e.g., dissipative magnetopheres---should be considered in modeling the radio polarization of the Crab pulsar's MP and IP. The nearly complete polarization of the LFC suggest that it originates at a different location and via a different mechanism than do the stronger MP and IP.  Finally, the fine time resolution and high signal-to-noise ratio in the MP data led to detection of statistically significant substructure in its pulse profile. We surmise that this substructure results from giant radio pulses occurring during the 144-minute observation.  {\\bf Acknowledgments}  The Westerbork Synthesis Radio Telescope (WSRT) is operated by ASTRON, the Netherlands Institute for Radio Astronomy, with support from NWO, the Netherlands Foundation for Scientific Research. AS acknowledges grant DEC-2011/03/D/ST9/00656 from the Polish National Science Centre; BWS, a Consolidated Grant from the UK Science and Technology Facilities Council; AKH, NASA grants Astrophysics Theory 12-ATP12-0169 and Fermi Guest Investigator 11-FERMI11-0052; AJvdH, Advanced Investigator Grant 247295 (PI: R.~A.~M.~J. Wijers) from the European Research Council; and SLO, RFE and MCW, support by NASA's Chandra Program.  \\appendix"
    },
    "1402/1402.5837_arXiv.txt": {
        "abstract": "The recently discovered X-ray line at about $3.5$~keV can be explained by sterile neutrino dark matter with mass, $m_s \\simeq 7$~keV, and the mixing, $\\sin^2 2\\theta \\sim 10^{-10}$. Such sterile neutrino is more long-lived than estimated based on the seesaw formula, which strongly suggests an extra flavor structure in the seesaw sector. We show that one can explain both the small mass and the longevity based on the split flavor mechanism where the breaking of flavor symmetry is tied to the breaking of the $B-L$ symmetry. In a supersymmetric case we find that the $7$~keV sterile neutrino implies the gravitino mass about $100$~TeV. ",
        "introduction": "Roughly a quarter of the Universe consists of dark matter. In spite of extensive dark matter searches conducted so far, its nature remains unknown. If dark matter is made of as-yet-unknown species of particles, they must be cold and very long-lived. The required longevity, however, does not guarantee the absolute stability of dark matter; it may decay into the standard model (SM) particles, enabling us to probe the nature of dark matter through indirect dark matter search.  Recently an unidentified X-ray line at about $3.5$~keV in the XMM-Newton X-ray observatory data of various galaxy clusters and the Andromeda galaxy was reported independently by two groups~\\cite{Bulbul:2014sua,Boyarsky:2014jta}. While there are a variety of systematic uncertainties that can affect the observed line energy and flux, it is interesting that such X-ray line can be explained by sterile neutrino dark matter~\\cite{Dolgov:2000ew,Boyarsky:2009ix,Kusenko:2009up,Abazajian:2012ys,Drewes:2013gca,Merle:2013gea}, a long-sought dark matter candidate, with mass about $7$~keV.  The sterile neutrino dark matter with mass in the keV range is known to decay radiatively into a photon and an active neutrino, producing a narrow X-ray line. The lifetime depends on its mass and the mixing with active neutrinos. For the mass of $7$~keV, the observed X-ray line can be explained by the mixing angle, $\\sin^2 2\\theta \\sim 7\\times 10^{-11}$~\\cite{Bulbul:2014sua,Boyarsky:2014jta}, which is just below the previously known X-ray bound. Such small mass and mixing can be partially understood by the split seesaw mechanism~\\cite{Kusenko:2010ik} or a simple Froggatt-Nielsen (FN) type flavor model~\\cite{Froggatt:1978nt}. In these scenarios, the seesaw formula~\\cite{seesaw} remains intact even in the presence of large mass hierarchy in the right-handed (sterile) neutrinos: \\bea \\left(m_\\nu\\right)_{\\alpha \\beta} = \\lambda_{I \\alpha}\\lambda_{I\\beta} \\frac{v^2 }{M_{I}}, \\label{seesaw} \\eea where $v \\equiv \\la H^0\\ra \\simeq174$~GeV is the vacuum expectation value (VEV) of the Higgs field, $\\lambda_{I \\alpha}$ denotes the Yukawa coupling of the right-handed neutrino $N_I$ with the lepton doublet $L_\\alpha$ and the Higgs field, and $M_I$ is the mass of the right-handed neutrino $N_I$. The point is that both the neutrino Yukawa coupling $\\lambda_{I \\alpha}$ and the right-handed neutrino mass $M_I$ are suppressed by either a geometrical factor or flavor charge in such a way that the suppression factors are cancelled out in the above seesaw formula. This is because the suppression mechanism is independent of the U(1)$_{B-L}$ breaking. The observed X-ray flux (as well as the previously known X-ray bound), however, requires that the sterile neutrino dark matter should be more long-lived than estimated based on the above seesaw formula.\\footnote{ It is known that, if the sterile neutrino comprises all the dark matter, its contribution to the light neutrino mass should be negligible to satisfy the X-ray bounds~\\cite{Asaka:2005an,Boyarsky:2006jm}. } The observed X-ray line therefore strongly suggests an extra flavor structure in the seesaw sector.  Before the discovery of the $3.5$~keV X-ray line, the present authors showed in Ref.~\\cite{Ishida:2013mva} that both the small mass and the small mixing just below the X-ray bound can be achieved in the split flavor mechanism; we introduce two $B-L$ Higgs fields, one of which is charged under single discrete flavor symmetry. The point is that the VEV of the $B-L$ Higgs leads to both breaking of the U(1)$_{B-L}$ symmetry and the flavor symmetry. In this letter we revisit the split flavor mechanism in light of the recent discovery of the unidentified X-ray line at $3.5$~keV, and show that the observed X-ray line can be nicely explained in the split flavor mechanism. In particular, we examine carefully the model parameters by taking account of numerical coefficients of order unity, while those numerical coefficients were set to be unity in Ref.~\\cite{Ishida:2013mva} for simplicity. In a supersymmetric case we will study the implications for supersymmetry breaking scale and show that the gravitino mass about $100$~TeV is favored by the observed X-ray line. ",
        "conclusions": "So far we have focused on the small mass and mixing of the sterile neutrino dark matter suggested by the recently discovered $3.5$~keV X-ray line, and have shown that these properties can be naturally realized in the split flavor mechanism. In order to explain the observations, we need to generate a right amount of the sterile neutrinos. Thermal production through mixings with active neutrinos, however, is inefficient for such small mixing angle, unless there is a large lepton asymmetry. Another possibility is through thermal production through U(1)$_{B-L}$ gauge interactions at high temperature~\\cite{Khalil:2008kp, Kusenko:2009up,Ishida:2013mva}; the abundance of the sterile neutrino is given by \\bea \\Omega_{\\rm DM} h^2 &\\sim& 0.2 \\lrfp{g_*}{106.75}{\\frac{3}{2}}  \\lrf{m_s}{7{\\rm\\,keV}}\\lrfp{M}{\\GEV{15}}{-4} \\lrfp{T_R}{5 \\times \\GEV{13}}{3}, \\eea where $T_R$ is the reheating temperature of the Universe after inflation. At such high reheating temperature, a right amount of the baryon asymmetry can be created through thermal leptogenesis~\\cite{Fukugita:1986hr} by the two heavy right-handed neutrinos~\\cite{Endoh:2002wm,Raidal:2002xf}. Non-thermal production may also work; see e.g. Ref.~\\cite{Petraki:2007gq}. Alternatively, if the $B-L$ symmetry is restored after inflation, the sterile neutrino $N_1$ will be thermalized through the $B-L$ gauge interactions. If there is a late-time entropy production of ${\\cal O}(10^2)$, its thermal abundance can be reduced so that it can explain the dark matter abundance~\\cite{Kusenko:2010ik}. In this case we need to introduce small explicit breaking of the discrete symmetry to make domain walls annihilate before dominating the Universe.  In this letter we have revisited the split flavor mechanism in light of the recent discovery of the X-ray line at $3.5$~keV in the XMM-Newton X-ray observatory data of various galaxy clusters and the Andromeda galaxy. In particular, the required small mixing angle, $\\sin^2 2\\theta \\sim 7\\times 10^{-11}$, implies that the sterile neutrino dark matter is more long-lived than estimated based on the seesaw formula, which strongly suggests an extra flavor structure in the seesaw sector. Note that the seesaw formula holds in the split seesaw mechanism or the simple FN flavor model. In the split flavor mechanism we introduce two $B-L$ Higgs fields, one of which is charged under discrete flavor symmetry. Most important, the breaking of flavor symmetry is tied to the breaking of the $B-L$ symmetry. We have shown in both non-supersymmetric and supersymmetric scenarios that the $7$~keV sterile neutrino with mixing $\\sin^2 2\\theta \\sim 7\\times 10^{-11}$ can be realized easily. The suppression of the mixing angle, namely, the longevity of the sterile neutrino dark matter with respect to the expectation based on the seesaw formula, is due to the mild hierarchy between the U(1)$_{B-L}$ breaking scale and the Planck scale. In the supersymmetric scenarios, the small mixing is partially due to the smallness of the gravitino mass compared to the Planck scale; the gravitino mass around $100$~TeV is favored by the observed X-ray line, which is consistent with the SM-like Higgs boson with $126$~GeV mass."
    },
    "1402/1402.0765_arXiv.txt": {
        "abstract": "Millisecond radio pulsars acquire their rapid rotation rates through mass and angular momentum transfer in a low-mass X-ray binary system. Recent studies of \\igr\\ and \\psr\\ have observationally demonstrated this link, and they have also shown that such systems can repeatedly transition back-and-forth between the radio millisecond pulsar and low-mass X-ray binary states. This also suggests that a fraction of such systems are not newly born radio millisecond pulsars but are rather suspended in a back-and-forth state switching phase, perhaps for giga-years. \\xss\\ has been previously suggested to be a low-mass X-ray binary, and until recently the only such system to be seen at MeV-GeV energies. We present radio, optical and X-ray observations that offer compelling evidence that \\xss\\ is a low-mass X-ray binary which transitioned to a radio millisecond pulsar state between 2012 November 14 and 2012 December 21.  Though radio pulsations remain to be detected, we use optical and X-ray photometry/spectroscopy to show that the system has undergone a sudden dimming and no longer shows evidence for an accretion disk. The optical observations constrain the orbital period to $6.913\\pm0.002$\\,hr. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) are systems in which a neutron star or black hole is orbited by a Roche-lobe-filling, main-sequence companion with a mass $\\lesssim 1$\\,\\MSun.  In systems with neutron star primaries the accretion disk can transfer both matter and angular momentum to the neutron star, thereby spinning it up and eventually producing a ``recycled\" millisecond radio pulsar (MSP; \\citealt{acrs82,rs82}).  The LMXB-MSP evolutionary link has been demonstrated by, e.g., (i) the observation of accretion-powered pulsations in \\sax\\ \\citep{wk98}; (ii) the discovery of the MSP \\psr, a system in which an accretion disk was previously observed \\citep{asr+09}; and (iii) the transition of \\igr/IGR\\,J18245$-$245 from an MSP to an accreting X-ray millisecond pulsar (AMXP) and back \\citep{pfb+13}.  More recently, \\psr\\ has re-entered a radio-quiet, accretion-disk state which shows much of the same X-ray and optical phenomenology observed in \\igr\\ \\citep{sah+13,tll+13,pah+14} --- though the system has not yet entered a state of full accretion onto the neutron star surface.  As such, we have witnessed a growing number of systems that straddle the classical definitions of LMXB and MSP, and which are helping us map this interesting evolutionary phase.  In addition to the aforementioned systems, recent radio pulsar searches have identified dozens of new ``black widow\" and ``redback\" systems \\citep{rob13}.  These are compact radio pulsar binary systems in which eclipsing of the pulsed signal is normally observed around superior conjunction of the neutron star, indicating that material is actively being ablated from the companion by the pulsar wind. The black widow systems are those in which the companion is very low mass ($\\sim 0.01$\\,\\MSun) and, likely degenerate.  The redbacks have likely non-degenerate companions that are significantly more massive (about 0.1 to 0.7\\,\\MSun).  The rapid growth in the known population of such sources has come from deep, targeted radio pulsation searches towards Galactic globular clusters (e.g.\\ \\citealt{rhs+05,hrs+07}) as well as towards unidentified $\\gamma$-ray sources found with the {\\it Fermi Gamma-ray Space Telescope} \\citep{rap+12}.  Since \\sax, \\psr\\ and \\igr\\ can all be classified as belonging to the black widow or redback families (though \\sax\\ has never been seen to pulse in radio it has shown evidence that a radio pulsar is active during the X-ray quiescent phase, see \\citealt{bsa+03} and \\citealt{cac+04}), it seems likely that in the coming decade some of these other, newly found MSPs will also be seen to transition to an X-ray active phase and back.  In fact, one may ponder whether these systems are gradually making a definitive transition from LMXB to MSP, or whether they are stuck transitioning back and forth for many giga-years (see \\citealt{ccth13} for evolutionary modeling of such systems).  \\psr's recent transition back to an LMXB-like state --- in which an accretion disk has returned \\citep{pah+14} --- has been accompanied by an additional observational surprise: namely that the MeV-GeV brightness of the system has quintupled in concert with the disappearance of the radio pulsar \\citep{sah+13}. The emission of $\\gamma$-rays is somewhat unexpected in LMXBs, and is more commonly seen in X-ray binaries with O or B-type companions in much wider orbits \\citep{dub13}. This is the first time that $\\gamma$-ray variability has been observed in an LMXB/MSP system, and it provides a new diagnostic tool while also raising new questions about the origin of this high-energy emission.  Prior to the increase of the \\psr\\ system's $\\gamma$-ray flux, the only LMXB thought to emit $\\gamma$-rays was \\xss. Discovered in the RXTE slew survey \\citep{sr04}, this system was initially classified as an $R=15.7$ cataclysmic variable by \\citet{mmp+06}, based on the presence of optical emission lines. Follow-up optical and X-ray observations cast doubt on this classification, suggesting instead that it is an LMXB \\citep{pre09,stei09,mfb+10}. Subsequently, \\citet{mfb+10} and \\citet{hsc+11} noted the positional coincidence of \\xss\\ with the \\textit{Fermi} $\\gamma$-ray source 1FGL\\,J1227.9$-$4852 (2FGL\\,J1227.7$-$4853). Based on that coincidence and similarities with \\psr, \\citet{hsc+11}, \\citet{mbf+13} and \\citet{ptl13} suggested that also \\xss\\ could harbour an active radio MSP.  Here we present radio, optical and X-ray observations to expand on our earlier report \\citep{bph+13} where we suggested that \\xss\\ has recently undergone a state transition similar to that observed in \\psr.  In the case of \\xss, however, it appears that the system has transitioned from an LMXB-like state to a radio MSP-like state, i.e.\\ the opposite switch that \\psr\\ has recently made \\citep{sah+13}. In \\S\\ref{Obs} we present our observations and analysis, in \\S\\ref{Res} our results, and end with a discussion in \\S\\ref{Disc}. ",
        "conclusions": "\\label{Disc}  Our radio, optical, and X-ray observations show that, sometime between 2012 November 14 and 2012 December 21, \\xss\\ transitioned to a new state where it is consistently and considerably fainter in these bands, and where the previous signs of an accretion disk (i.e.\\ double-peaked optical emission lines) have disappeared. Subsequent to our report of this transition in \\citet{bph+13}, \\citet{tkl13} reported a similar decrease in brightness of \\xss\\ in $\\gamma$-rays, while \\citet{cmm+14} confirm the absence of emission lines at all orbital phases.  This phenomenology shows many parallels with \\psr, which has recently transitioned from being a radio millisecond pulsar to a state where it resembles a quiescent LMXB \\citep{asr+09,wat+09,sah+13,tll+13,pah+14}. Compared with \\psr's recent behaviour, however, the \\xss\\ observations presented here strongly suggest the reverse transition, i.e.\\ from an LMXB-like state in 2012 to one where the accretion disk is absent in 2013. Besides the presence of an accretion disk during the LMXB state, the \\xss\\ X-ray light-curve presented by \\citet{mbf+13} displays frequent low flux states in brightness, similar to what is seen in the recent LMXB state of \\psr\\ \\citep{pah+14} as well as during \\igr's X-ray active phase \\citep{pfb+13}. The cause of the low X-ray flux states in both systems is presently unknown. The sinusoidal $U$-band light-curve that is seen in the present state of \\xss\\ is comparable to the light-curves presented in \\citet{wwp04} and \\citet{ta05} and can be understood as irradiation of the companion by the neutron star.  A radio source coincident with \\xss\\ was detected by \\citet{hsc+11} in ATCA observations obtained in 2009 when the source was in the LMXB state. Our follow-up ATCA observations obtained after the transition show it has decreased in brightness by at least a factor of 6. Recent Very Large Array observations of \\psr\\ after it returned to the LMXB state show a flat spectrum radio source (Deller et al.\\ in prep.). All these parallels between \\xss\\ and \\psr\\ lead us to conclude that \\xss\\ harbours an active rotation-powered millisecond pulsar (see also \\citealt{ptl13}) and, based on the present information, one could classify \\xss\\ as a ``redback'' system in its current state.  Thus far, however, we have failed to detect radio pulsations from \\xss.  Though such a detection would show without a doubt that the system has transitioned from an LMXB state to an MSP state, the lack of radio detection in no way rules out such a state transition.  Given the high inferred mass ratio, $q=0.53$ \\citep{mcm+13}, it appears that \\xss\\ would be observed as a redback in its present state. In analogy with other redback systems, we expect the pulsar to be eclipsed for $\\sim 50$\\% of the time at $\\sim 1.4$\\,GHz, and it is possible that the system is enshrouded a much larger fraction of the time.  Indeed, the recent examples of PSR\\,J2339$-$0533 \\citep{rs11} and PSR\\,J1311$-$3430 \\citep{rrc+13} show that the radio detectability of such pulsars may be exceedingly poor, even though the radio pulsar mechanism is clearly active. For this reason, continued radio pulsar searches may eventually be successful in detecting the source. It is also simply possible that the radio pulsar is too weak to detect with the current Parkes data, too highly accelerated without using a far more complex acceleration search, or not beamed towards the Earth.  \\psr's transition from MSP to LMXB was accompanied by a 5-fold increase in $\\gamma$-ray flux \\citep{sah+13}. Assuming that \\xss\\ is mirroring such a transition, we would thus expect a marked decrease in $\\gamma$-ray brightness. Using the optically and X-ray-derived epoch of the state change, \\citet{tkl13} report that the $\\gamma$-ray brightness has indeed decreased by a factor of 1.5 to 2, but note that the decrease appears more gradual than what was seen in \\psr. Indeed, the $\\gamma$-ray light-curve of \\xss\\ appears more complicated than that of \\psr\\, and we are investigating whether there is contamination from another source.  \\psr's transition from radio MSP to LMXB is constrained to have happened within a 2-week window \\citep{sah+13}, and may well have happened even more abruptly.  \\xss's mirror transition from LMXB to an MSP-like state is constrained to a 5-week period, and is thus also quite rapid.  An additional comparison is provided by \\igr's 2013 transition from accreting X-ray MSP to an observable radio MSP, within a period of less than 3 weeks.  It thus appears that the back-and-forth transitioning of such systems between MSP and LMXB is a rapid process which high-cadence radio and X-ray monitoring can constrain further.  In particular, though practically difficult to achieve, near-daily, joint radio and X-ray monitoring would either detect or strongly constrain whether there is a lag between the X-ray brightening and disappearance of the radio MSP.  Besides the similarities with \\psr, there are also differences. \\xss's high reported mass ratio $q = 0.53$ \\citep{mcm+13} suggests a significantly more massive companion compared with \\psr\\ ($M_\\mathrm{c}=0.2$\\,\\MSun) and the majority of the other observed redbacks (see \\citealt{rob13} for a review). Such a high companion mass is not unprecedented, however; PSR\\,J1723$-$2837 \\citep{cls+13} has an inferred $M_\\mathrm{c,min} = 0.7$\\,\\MSun\\ companion and PSR\\,J2129$-$0428 (Hessels et al., in prep.) has a $M_\\mathrm{c,min} = 0.5$\\,\\MSun\\ companion. A much larger number of LMXB/MSP transition systems will have to be characterized before it is possible to determine the role that companion mass has on the frequency of state switches.  Thus far we have observed similar back-and-forth transitioning from three redbacks; \\psr, \\igr, and now also in \\xss.  Should we expect the same from the black widows?  For the accreting X-ray millisecond pulsars, roughly half have black-widow-like companions and half have redback-like companions \\citep{pw12}. There is no obvious correlation between the recurrence timescales in such system and the companion mass.  Why have such systems not been observed as radio pulsars? Possibly simply because they are mostly very distant, which is problematic both because radio MSPs are intrinsically quite faint and because distant sources at low Galactic latitudes will be scattered.  The observations and analysis presented in this paper show that \\xss\\ has undergone a state transition similar to those observed in \\psr\\ and \\igr. The analysis of recently obtained \\textit{XMM/Newton} and \\textit{Chandra} X-ray observations of \\xss\\ may shed further light on the presence of an active radio MSP, though multiwavelength observations are certainly warranted. For example, the present absence of the accretion disk allows for optical spectroscopy of the main sequence companion. This will provide constraints on the masses of the companion and the neutron star. Finally, regular optical and X-ray will help identify future transitions, constraining the timescales involved in the transitions and possibly trigger detailed investigations during the transitions."
    },
    "1402/1402.2283_arXiv.txt": {
        "abstract": "We study the minimal ``bathtub\" toy model as an analytic tool for capturing key processes of galaxy evolution and identifying robust successes and challenges in reproducing observations at high redshift.  The source and sink terms of the continuity equations for gas and stars are expressed in simple terms from first principles.  The assumed dependence of star-formation rate (SFR) on gas mass self-regulates the system into a unique asymptotic behavior, which is approximated by an analytic quasi-steady-state solution (QSS).  We address the validity of the QSS at different epochs independent of earlier conditions.  At high $z$, where the accretion is assumed to consist of gas only, the specific SFR is robustly predicted to be sSFR$\\,\\simeq[(1+z)/3]^{5/2}\\Gyr^{-1}$, slightly higher than the cosmological specific accretion rate, in agreement with observations at $z=3\\!-\\!8$.  The gas fraction is expected to decline slowly, and the observations constrain the SFR efficiency per dynamical time to $\\epsilon\\!\\simeq\\!0.02$. The stellar-to-virial mass ratio $\\stv$ is predicted to be constant in time, and the observed value requires an outflow mass-lading factor of $\\eta\\!\\simeq\\!1\\!-\\!3$, depending on the penetration efficiency of gas into the galaxy.  However, at $z\\!\\sim\\!2$, where stars are also accreted through mergers, the simplest model has an apparent difficulty in matching observations.  The model that maximizes the sSFR, with the outflows fully recycled, falls short by a factor $\\sim\\!3$ in sSFR, and overestimates $\\stv$. With strong outflows, the model can reproduce the observed $\\stv$ but at the expense of underestimating the sSFR by an order of magnitude.  We discuss potential remedies including a bias due to the exclusion of quenched galaxies. ",
        "introduction": "\\label{sec:intro}  The bathtub toy model is simply the equation of conservation of gas mass in a galaxy, or in one of its components, sometimes combined with the analogous equation for stellar mass. The gas equation describes the net change of gas mass in the inter-stellar medium (ISM) as a difference between source terms and sink terms. When a whole galaxy is considered, the source term is the accretion rate, dictated by cosmology, but it can include recycled gas returning to the galaxy. The main sink terms are the rate at which gas turns into stellar mass and the gas outflow rate. The key for making the solution of this equation converge to a unique solution independent of the initial conditions is that the SFR is assumed to be proportional to the gas mass, and the other sink terms, especially the outflow rate, are assumed to be proportional to the SFR. This generates an interplay between the two components, where more gas mass allows a higher SFR which in turn reduces the gas mass, thus driving the system into a self-regulated state that is determined by the relative efficiencies of accretion, SFR and outflows. When the accretion rate and the SFR timescale vary slowly enough, the asymptotic solution can be approximated by a quasi-steady state (QSS, sometimes termed ``equilibrium\") that can be derived analytically.  \\smallskip The bathtub model is useful because it captures the key processes in a very simple and transparent way that makes it easy to trace their roles in the global evolution. It is appealing because it converges to a unique attractor solution, and even more so because of the analytic QSS solution. One of its attractive features is that at any time this solution is unique, independent of the initial conditions and independent of whether the same model was valid at earlier epochs. The model is thus useful even if its ingredients with specific values of the model parameters are valid only in a relatively short cosmological time interval.  \\smallskip The bathtub model has become a very useful tool in understanding in simple terms the gross features of galaxy evolution. As summarized in \\citet{dekel13}, it has been applied in different ways to study different aspects of the the evolution of a whole galaxy \\citep{bouche10,dave11I,dave12,kd12,lilly13,birrer_lilly14}. Alternatively, it has been applied to the evolution of discs undergoing violent disk instability (VDI) where another sink term is the mass inflow from the disc to the central bulge \\citep{dsc09,cdg12,gdc12,forbes13}. Recently it has been applied to the mass evolution of individual giant clumps in VDI disks \\citep{dekel_mass_14}.  \\smallskip The bathtub toy model is not a replacement for the more elaborate semi-analytic models (SAMs) or for full hydro-cosmological simulations. This is a toy model, where on one hand the physical processes are represented by very idealized recipes but on the other hand their effects are fully transparent. It can thus serve for solid qualitative constraints on the parameters that characterize the accretion, SFR and outflows. It also serves for revealing robust successes or difficulties in understanding the origin of observed galaxy properties.  \\smallskip In the first part of this paper, we present the bathtub toy model and its ingredients in the case of a galaxy accreting from the cosmic web. We express the source and sink terms using the simplest possible scaling relations that we try to motivate from first principles rather than from simulations or observations. These simplified approximate expressions, in the Einstein-deSitter cosmological regime (approximately valid at $z>1$), allow an analytic derivation of the QSS solution. We numerically compute the asymptotic solution, derive the analytic QSS approximation, analyze its range of validity, and reveal the role played by each process in major observable quantities at different cosmological epochs.  \\smallskip In the second part we address three rather intriguing observational results at high redshift, which we describe in more detail in \\se{comp_obs}. First is the average specific star-formation rate (sSFR) for massive star-forming galaxies (SFGs) as a function of redshift, which current estimates show to decline from $\\sim\\! 10 \\Gyr^{-1}$ at $z \\!\\sim\\! 8$ through $\\sim\\! 2 \\Gyr^{-1}$ at $z \\!\\sim\\! 2$ to $\\sim\\! 0.1-0.2 \\Gyr^{-1}$ at $z=0$ \\citep{whitaker12_ssfr,stark13,labbe13,salmon14}. Second is the gas fraction in the galaxy, $\\fg$, deduced to be $\\sim\\! 0.5$ at $z \\!\\sim\\! 2$ and rather slowly declining with time at high redshifts \\citep{daddi10,tacconi13}. Third is the stellar-to-virial mass ratio, $\\stv$, estimated to be independent of redshift at $z=0-4$, with a peak of $\\sim\\! 0.025$ for halo masses in the range $(0.5-5)\\times 10^{12}\\msun$ \\citep[e.g.][see \\se{comp_obs}]{moster10,moster13,behroozi13}. Given the universal baryon fraction $\\fb \\simeq 0.17$, this implies that only $\\sim\\! 0.15$ of the baryons that were supposed to enter the halo have made it to the central galaxy. The sSFR that exceeds the expected specific accretion rate at all times, the high gas fraction as late as $z \\!\\sim\\! 2$, and the very low $\\stv$, all pose non-trivial theoretical challenges.  \\smallskip In addition, the observed Kennicutt-Schmidt star-formation law \\citep{kennicutt89}, which can be interpreted as a relation between the SFR in a galaxy and its gas content, reveals a rather constant SFR efficiency in a free-fall time at the level of one to a few percent \\citep{krum_tan07,krum_thomp07, daddi10_sfr,genzel10,kdm12}. Then, significant outflows are observed in high-redshift galaxies \\citep{steidel10,genzel11}, with mass-loading factors ranging from below unity to 10. In both cases the uncertainties are large, so we prefer to treat these as free variables.  \\smallskip The paper is organized as follows. In \\se{model} we present the bathtub model and its ingredients. In \\se{qss} we describe the quasi-steady-state solution and work out analytically observable quantities of interest. In \\se{validity} we study the validity of the QSS solution. In \\se{comp_obs} we make a simple comparison of the model to observations. In \\se{disc} we discuss the successes and failures of the model in reproducing the observations, and address the assumptions made in the toy model. In \\se{conc} we conclude our findings concerning the toy model and the comparison to observations. ",
        "conclusions": "\\label{sec:conc}  The minimal bathtub toy model is shown to be a useful tool for tracing the roles played by key processes of galaxy evolution and identifying major successes and tentative failures in reproducing observations at high redshift. This model is based on simplified but robust continuity equations for the gas and for the stars. Because of the generic monotonic dependence of the SFR (and therefore outflow rate) on gas mass, the system is self-regulated into a unique asymptotic behavior, which can be approximated by a quasi-steady-state solution. The simple time and mass dependence of the average accretion rate into galaxies allows an approximate analytic quasi-steady state solution.  \\smallskip We derived the analytic QSS solution and evaluated the associated deviation from the exact solution for three observables: the sSFR, the gas fraction, and the stellar-to-virial mass fraction. We showed that the error is negligible for the first two, and is limited to the level of tens of percent for the third quantity. We studied the range of validity of the QSS solution, and found that the errors are smaller at later epochs, and when the outflows are stronger.  \\smallskip In the QSS regime, the average sSFR is proportional to the specific accretion rate, sAR, with an absolute value that is insensitive to the SFR efficiency in a dynamical time $\\epsilon$ and to the outflow mass-loading factor $\\eta$. The gas fraction is determined by $\\epsilon$, and $\\stv$ is driven by $\\eta$.  \\smallskip At high redshifts, where the accretion is assumed to consist of gas only, the simple toy model reproduces the observations rather well with no need for fine tuning. The specific SFR is predicted with no free parameters to be sSFR$\\,\\simeq [(1+z)/3]^{5/2} \\Gyr^{-1}$, slightly higher than the cosmological sAR, in general agreement with the rather noisy observed sSFR at $z=3-8$. The observed gas fraction constrains the SFR efficiency in a dynamical time to $\\epsilon \\simeq 0.02$. The low $\\stv$ indicated from observations requires an outflow mass-loading factor of $\\eta \\simeq 1-3$, for a penetration efficiency of fresh gas into the galaxy of $\\p = 0.5-1$ respectively. Thus, the main features of galaxy evolution at high redshifts are captured by the simplest toy model.  \\smallskip However, at $z \\!\\sim\\! 2$, where stars are also accreted, through mergers, there is a difficulty in matching the observations. The model with the highest possible sSFR, where the outflows are fully recycled, falls short by a factor $\\sim\\! 3$ compared to the observed sSFR, and it overestimates $\\stv$. With $\\eta \\!\\sim\\! 3$, the model reproduces the latter but underestimates the sSFR by an order of magnitude. Thus, the toy model points at a robust challenge to theory at $z \\!\\sim\\! 2$.  \\smallskip A potential way to ease the tension at $z\\!\\sim\\!2$ is by massive recycling, as in the case modeled by $\\eta =0$ in the minimal bathtub model. A more sophisticated implementation of recycling in the model may be required for a better match. The missing population of quenched galaxies in the observed average sSFR hints at an additional promising remedy. The galaxies where $\\eta$ or $\\fsa$ are higher than average drive the predicted average sSFR down, and the same galaxies are indeed likely to be quenched in reality. Once these are removed, the predicted sSFR would become higher by a factor of 2-3 and thus closer to the observed value. Finally, the discrepancy at $z \\!\\sim\\!2$ may be due to a breakdown of the QSS solution, e.g., by strong fluctuations in the accretion rate. Alternatively, one could not deny the possibility of some fundamental flaw in the assumptions that lie at the basis of the bathtub toy model. Identifying such a flaw would be very interesting.  \\smallskip One should conclude with a reminder that the bathtub toy model, at least in its trivial form discussed here, is not meant to provide a detailed model that perfectly matches the observations. It is useful for a simple and crude qualitative study of the central elements of galaxy evolution, and for pointing out robust successes or major potential failures in reproducing certain observations, which should trigger a more detailed study of these issues. This has been demonstrated here by showing the robust match to observations at $z=3-8$, and by pointing out the non-trivial discrepancy at $z\\!\\sim\\!2$.  \\smallskip On the other hand, the bathtub toy model can be extended to address other major observables, such as the evolution of gas metallicity in and outside galaxies and the scaling relations involving metallicity, as well as the hot gas content in the circum-galactic medium filling the halo and the associated with hot-mode accretion into the galaxy. Such extensions can be found, e.g., in \\citet{kd12}, \\citet{dave12}, and \\citet{lilly13}, and they are shown to be good approximations (Pipino et al.~2014, in preparation)."
    },
    "1402/1402.5140_arXiv.txt": {
        "abstract": "Recently, \\citeauthor{DiazTrigo13} reported an \\textit{XMM-Newton} detection of relativistically Doppler-shifted emission lines associated with steep-spectrum radio emission in the stellar-mass black hole candidate 4U 1630--47 during its 2012 outburst. They interpreted these lines as indicative of a baryonic jet launched by the accretion disk.  Here we present a search for the same lines earlier in the same outburst using high-resolution X-ray spectra from the \\textit{Chandra} High-Energy Transmission Grating Spectrometer. While our observations (eight months prior to the \\textit{XMM-Newton} campaign) also coincide with detections of steep spectrum radio emission by the Australia Telescope Compact Array, we find no evidence for any relativistic X-ray emission lines. Indeed, despite $\\sim5\\times$ brighter radio emission, our \\textit{Chandra} spectra allow us to place an upper limit on the flux in the blueshifted Fe\\,{\\sc xxvi} line that is $\\gtrsim20\\times$ weaker than the line observed by \\citeauthor{DiazTrigo13} We explore several scenarios that could explain our differing results, including variations in the geometry of the jet or a mass-loading process or jet baryon content that evolves with the accretion state of the black hole. We also consider the possibility that the radio emission arises in an interaction between a jet and the nearby ISM, in which case the X-ray emission lines might be unrelated to the radio emission. ",
        "introduction": "\\label{sec:intro} \\defcitealias{DiazTrigo13}{DT13}  Collimated jets are ubiquitous in accreting systems across the mass scale, from protostars (e.g., \\citealt{CarrascoGonzalez10}) to supermassive black holes in AGN at the centers of galaxy clusters (e.g., \\citealt{McNamara00,Fabian00}). But the omnipresence of jets in astrophysical systems raises important questions. Are all these jets produced by the same processes? Since one of the canonical jet formation mechanisms (the Blandford-Znajek, or BZ, mechanism; \\citealt{BlandfordZnajek,Tchekhovskoy11}), involves the rotation power of a black hole, it certainly cannot operate in all systems. An alternative is the Blandford-Payne mechanism (BP; \\citealt{BlandfordPayne}), in which the requisite rotation power comes from a rotating accretion disk. As a consequence, one expects BP jets to be baryonic, since they are magnetocentrifugal outflows from a gaseous reservoir. BZ jets, on the other hand, might be electromagnetically dominated.  In this context, it is particularly interesting that accreting stellar-mass black holes appear to produce two different kinds of jets (steady and transient; for reviews see \\citealt{FBG04,Fender09}). Observationally, it remains unclear whether these two types of jets correspond to the two mechanisms described above; a suggested association between the transient jets and the BZ mechanism has lately been the subject of significant debate (\\citealt{Fender10,Narayan12,Steiner13,Russell13,King13b}). The formation of steady jets is no less significant, as it has been suggested (\\citealt{NL09}; see also \\citealt{M08,N11a,N12b,King13} and references therein) that ionized accretion disk winds may be able to quench steady jets. Different processes may launch astrophysical jets at different scales or under different conditions, but because a jet's energy requirements are sensitive to its baryon content (see, e.g., \\citealt{F99,Gallo05,Heinz06,Punsly13}), these considerations have broad and concrete implications for radio-mode feedback from accreting systems.  \\begin{figure*} \\centerline{\\includegraphics[width=\\textwidth]{f1}} \\caption{Monitoring of the 2012 outburst of 4U 1630--47. (a) 2--20 keV MAXI light curve. (b) MAXI hardness ratio (10--20/2--4 keV). (c) 15--50 keV \\textit{Swift}/BAT light curve. (d) 5.5 GHz and 9 GHz radio fluxes as measured by ATCA by us and \\citetalias{DiazTrigo13}. Dashed and dash-dotted vertical lines indicate the times of the \\textit{Chandra} and \\textit{XMM} observations, respectively. (e) MAXI hardness-intensity diagram. The colors/symbols are rough indicators of the accretion state.\\label{fig:lc}} \\end{figure*}  Recently, \\citeauthor{DiazTrigo13} (\\citeyear{DiazTrigo13}; hereafter \\citetalias{DiazTrigo13}) reported a detection of relativistically-blueshifted emission lines associated with radio emission in the stellar-mass black hole candidate 4U 1630--47, a well-known recurrent transient  (\\citealt{Jones76,Kuulkers97b,Tomsick05}) that went into outburst in December 2011 (\\citealt{N13a}). \\citetalias{DiazTrigo13} reported on two observations of the source with \\textit{XMM-Newton} and the Australia Telescope Compact Array (ATCA) in 2012 September. In their first observation (September 10-12), they found a featureless X-ray spectrum and no radio emission. In their second observation (September 28-29), however, they discovered optically thin radio emission and X-ray emission lines at energies that do not correspond to strong lines from abundant elements. They argued that the emission lines originated in a relativistic jet with a speed of $\\sim0.66c,$ where $c$ is the speed of light. In their interpretation, the lines corresponded to red- and blueshifted Fe\\,{\\sc xxvi} Ly$\\alpha$ and blueshifted Ni\\,{\\sc xxvii} He$\\alpha$.  In this paper, we report on ATCA and \\textit{Chandra} High-Energy Transmission Grating Spectrometer (HETGS; \\citealt{C05}) observations of the same outburst of 4U 1630--47. Although we detect radio emission at levels brighter than those reported in \\citetalias{DiazTrigo13}, we do not detect the X-ray emission lines seen later in the outburst. Indeed, our \\textit{Chandra} spectra place tight upper limits on such features, and we can conclusively say that radio emission is not universally associated with relativistically Doppler-shifted emission lines in 4U 1630--47. While our results do not and cannot directly address the argument for the presence of baryons in the 2012 September jet in 4U 1630--47, we highlight the necessity of a complex relationship between X-ray and radio emission, considering not only the possibility of undetected baryons in the January and June jets or baryonic jets confined to a particular accretion state, but also an alternative scenario in which the radio emission arises in a jet-ISM interaction and could be decoupled from the X-ray emission. ",
        "conclusions": "\\label{sec:concl} The baryon content of astrophysical jets has important consequences for our understanding of their formation and their effects on their environment, but it is difficult to measure directly. \\citet{DiazTrigo13} have argued that their detection of relativistically Doppler-shifted emission lines in 4U 1630--47 in September 2012 (MJD 56181) implies the presence of a baryonic jet, most likely launched from the rotating accretion disk (\\citealt{BlandfordPayne}). We cannot directly test this interpretation, but in this paper we have used \\textit{Chandra} HETGS observations from earlier in the same outburst (January/June, MJD $\\sim55954/56081$) to demonstrate that radio emission in 4U 1630--47 is not always associated with relativistic X-ray emission lines. The unique behavior observed with \\textit{XMM-Newton} and ATCA later in the outburst may be a special case, dependent on additional processes in the accretion flow around the black hole. However, because much of the observed radio emission from 4U 1630--47 in 2012 is consistent with jet-ISM interactions, we cannot rule out scenarios where the X-ray emission lines originate in a hot outflow that is not physically related to the radio source. Further analysis of the X-ray and radio emission (Neilsen et al., in preparation) may help test our interpretation of this exciting physics.\\vspace{-2mm}"
    },
    "1402/1402.0848_arXiv.txt": {
        "abstract": "In this short communication we present a simple internal mechanism that accounts for the recently observed anti-glitch in magnetar 1E 2259+586.  We propose that the decay of an internal toroidal magnetic field component would de-stabilize an originally stable prolate star configuration. Then, the subsequent rearrangement of the stellar structure would give rise to a ``more spherical'' configuration, resulting in a sudden spin-down of the whole star.  We present here some order of magnitude calculation to give confidence to this scenario, using the simplest analytical stellar model and let more detailed calculations for a more technical future paper. ",
        "introduction": "Magnetars are neutron stars powered by their strong internal magnetic fields \\citep{duncan1992}. The detailed study of the temporal behaviour of magnetar's emission can be used to constrain both their external dipolar magnetic field and internal structure \\citep{kouve1998,glitches}. Despite the fact that all neutron stars suffer a long term spin-down, due to their dipole magnetic field decay, many sudden spin-up, known as glitches in the literature, have been observed in pulsars and magnetars \\citep{glitches1,glitches2}. Recently, clear evidence of the first sudden spin-down detection was observed in magnetar 1E 2259+586 \\citep{antiglinch}, a so-called ``anti-glitch''. There, the authors propose two different interpretations for the observational evidence: (i) an anti-glitch event followed by a normal glitch (after $\\sim$100 days) and (ii) a sequence of two anti-glitches (separated by $\\sim$50 days). Spin-down events were also observed in magnetars SGR 1900+14 \\citep{sgr1900} and 4U 0142+61 \\citep{4u}, and in the high magnetic field pulsar PSR J1846--0258 \\citep{psrj}, but none of them is considered as an anti-glitch, because of their long timescales ($17-127$~days).  The Anomalous X-ray Pulsar 1E 2259+586, with a $\\sim 7$~s period, is a magnetar with a characteristic age $\\sim 10^4$~yr and a spin-inferred surface dipolar magnetic field of $B_{d} \\sim 5.9 \\times 10^{13}$~G, which poses a minimum value for the internal magnetic field strength. 1E 2259+586 has been monitored with {\\it Rossi X-ray Timing Explorer} and {\\it Swift X-ray Telescope} over almost the last two decades, showing a stable spin-down rate, with the exception of two spin-up glitches in 2002 \\citep{kaspi2003} and 2007 \\citep{idem2012}, a timing event in 2009 \\citep{idem2012}, and this anti-glitch in 2012.  From \\cite{antiglinch}, it follows that the observational fact that has to be explained by any model is a total change in frequency in the complete $\\sim 100$ day event was $\\sim -5 \\times 10^{-7}$~Hz, that can be interpreted as two instantaneous changes in spin frequency $(|\\Delta \\nu / \\nu| \\gtrsim  10^{-7})$. In order to adjust the timing data from {\\it Swift}, the authors propose two models: (i) an anti-glitch in which $\\Delta \\nu / \\nu = -3.1(4) \\times 10^{-7}$ followed by a spin-up event, of amplitude $\\Delta \\nu / \\nu = 2.6(5) \\times 10^{-7}$; (ii) an anti-glitch in which $\\Delta \\nu / \\nu = -6.3(7) \\times 10^{-7}$ followed by a second anti-glitch in which $\\Delta \\nu / \\nu = -4.8(5) \\times 10^{-7}$. Based on a bayesian analysis, model (ii) is favoured \\citep{hu2013}.  Several explanations for this anti-glitch event have been proposed, based both on external \\cite{antiglinch-explanations} and internal \\citep{duncan2013} origins. Despite searches in radio and X-ray wavelengths, no surrounding afterglow was detected \\citep{antiglinch}, arguing against a sudden particle outflow or wind-driven scenario. In this sense, a most promising approach seems to be a an internal rearrangement of the star. With this short communication we aim to provide a simple mechanism where the sudden spin-down of the star occurs while the stellar structure reaccommodates after a substantial decay of the internal toroidal magnetic field component. A similar scenario was also suggested to account for the SGR 1900+14 event \\citep{ioka}.  In \\citep{cutler.2002,mag-deform}, the authors developed a formalism to model magnetically deformed neutron stars. In their work, deformations are calculated for a uniform density star with a mixed poloidal-toroidal magnetic field configuration, rotating with angular velocity, which, in the case of magnetars, as 1E 2259+586, is completely irrelevant, because of their long periods. The main result of their work is that while for strong poloidal magnetic fields, as for rapid rotation, the stars tend to oblate to keep mechanically stable, if internal toroidal fields dominate the magnetic field configuration, prolate stars are favoured.  For volume preserving $l=2$ mode, and mixed toroidal-poloidal magnetic field, the quadrupolar distortion for incompressible stars of uniform density is given by:  \\begin{equation} \\label{epsilon} \\epsilon = \\frac{I_{ zz} - I_{xx}}{I_{zz}} = - \\frac{25R^4}{24G_NM^2}\\left(\\left<B_t^2\\right>-\\frac{21}{10}\\left<B_p^2\\right>\\right), \\end{equation}  \\noindent where $\\left<B_t^2\\right>$ is the mean value of the square of the toroidal magnetic field strength, $\\left<B_p^2\\right>$ is the mean value of the poloidal component, $R$ the radius of the undeformed star, $G_N$ the gravitational constant and $M$ the mass of the neutron star.  Despite that a purely toroidal magnetic field is unstable \\cite{instab-toro, instab-polo1, instab-polo2, stable}, an additional poloidal component with energy $E_p/E_t = B_p^2/B_t^2 \\sim 1-5\\%$ stabilizes the magnetic field configuration \\cite{reisenegger2013}, allowing us to neglect the poloidal contribution to Eq. \\ref{epsilon} as $\\left<B_p^2\\right> \\ll \\left<B_t^2\\right>$. ",
        "conclusions": "A very simple model to explain the anti-glitch observed in magnetar 1E 2259+586 is presented here. We propose that a natural decay of the toroidal magnetic field component, of approximately $\\sim 10\\%$, from an initial $\\left<B_t^i\\right> \\gtrsim 10^{15}$~G, would be enough to de-stabilize an originally prolate configuration for the stellar structure, reducing it to a ``more spherical'' one. As a result, a sudden change in the moments of inertia of the star produces a net spin-down, as the one observed in \\cite{antiglinch}. Detailed studies of the magnetic field evolution in neutron stars show that a magnetic field decay of $\\sim 10\\%$ as the one needed in our model is easily achieved after $t \\lesssim 10^4$~yr for a magnetar like 1E 2259+586 \\citep{vigano2013}.  It is also worth to notice that gravitational waves potentially observed by advanced LIGO, should be emitted under this scenario due to oscillations induced by the changes in the stellar structure.  More involved calculations using more realistic neutron star configurations are being carried away \\cite{con-fede-prep} to refine these results."
    },
    "1402/1402.2518.txt": {
        "abstract": "Minerals detection over large volume of spectra is the challenge addressed by current hyperspectral imaging spectrometer in Planetary Science. Instruments such OMEGA (Mars Express), CRISM (Mars Reconnaissance Orbiter), M$^{3}$ (Chandrayaan-1), VIRTIS (Rosetta) and many more, have been producing very large datasets since one decade. We propose here a fast supervised detection algorithm called LinMin, in the framework of linear unmixing, with innovative arrangement in order to treat non-linear cases due to radiative transfer in both atmosphere and surface. We use reference laboratory and synthetic spectral library. Additional spectra are used in order to mimic the effect of Martian aerosols, grain size, and observation geometry discrepancies between reference and observed spectra. The proposed algorithm estimates the uncertainty on ``mixing coefficient'' from the uncertainty of observed spectra. Both numerical and observational tests validate the approach. Fast parallel implementation of the best algorithm (IPLS) on Graphics Processing Units (GPU) allows to significantly reduce the computation cost by a factor of \\textasciitilde{}40. ",
        "introduction": "Various methods have been proposed to detect surface chemical species (minerals, ice) on large dataset of hyperspectral images. Supervised methods (knowing the spectra of the chemical species you want to detect) are widely used, for instance : band ratio techniques (\\citep{Poulet_OMEGA-MnxMap_JGR2007,Ehlmann_PhylosolicateMars_Nature2011,Carter_Automatedprocessingplanetary_PaSS2013}), linear unmixing (\\citep{Combe_MELSUM_PSS2008,Themelis_OMEGAhyperspectral_PSS2012}), wavelet based detection (\\citep{Schmidt_Wavanglet_IEEETGRS2007,Gendrin_ClassifHyperspectral_JGR2006}), correlation based detection - Spectral Angle Mapper (\\citep{Kruse1993}). Other linear non-supervised techniques (estimating the spectra directly from the scene) have been proposed, using Independent Component Analysis ICA (\\citep{2005LPI....36.1623F,Erard_ICAvirtis_JGR2009}) or Bayesian methods under constrained (\\citep{Moussaoui_JADE-BPSS_Neurocomp2008,Schmidt_ImplementationBPSS_TGRS2010}). Also some algorithms only extract the endmember spectra such: Pixel Purity Index (\\citep{1995Boardman}), N-FINDR (\\citep{Winter_NFINDR:algorithm_1999}), or graph-based segmentation (\\citep{Gilmore_SuperpixelHyperspectral_JGR2011}) .  The linear mixture is still valid in non-linear intimate mixture, but with significant difference between retrieved ``mixing coefficient'' and actual abundance (\\citep{Mustard_Photometry_JGR1994,Mustard_nonlinear_JGR1998}). Nevertheless, linear unmixing is satisfactory for mineral detection (\\citep{Combe_MELSUM_PSS2008,Themelis_OMEGAhyperspectral_PSS2012}). It has the significant advantage to deal with complex mixture of a large variety of candidate minerals (see fig. \\ref{fig:32ReferenceSpectra}), whereas band ratio methods fails, due to the lack of defined reference ``continuum'' wavelength channel. In such situation each wavelength may sample an absorption band of a particular mineral and no reference wavelength can be found. If band ratio's are adapted to detect minerals over flat spectra, methods using the entire wavelength channels (such linear unmixing) are adapted to detect minerals over a complex variety of background. Specifically, we draw attention that the band ratio's of the CRISM summary products are only relevant in case of a pure mineral detection but may be irrelevant in case of a mixture. We propose here to address the challenge to detect one mineral type alone or a mineral type in an assemblage, from a very diverse potential endmember spectra dataset.  Linear unmixing is relatively simple enough to provide a fast implementation. Furthermore supervised linear unmixing has the advantage that its interpretation is directly and automatically provided in terms of mineralogical class on the bulk hyperspectral image. Dimensionality reduction methods, such unsupervised algorithms or endmember extractions are not required in linear unmixing. Also linear unmixing does not require manual interpretation and identification of spectra.  Previous linear unmixing algorithms have proposed to optimize the reference spectral library in order to solve positivity (\\citep{Combe_MELSUM_PSS2008}), forcing for sum-to-one constraints on the mixing coefficients (\\citep{Roberts_MappingChaparralin_RSoE1998}). \\citep{Roberts_MappingChaparralin_RSoE1998} discusses the benefits of constraining the number of endmembers from the spectral library. They have shown that the solutions with a mixture of two endmembers have less overlap that three endmembers mixtures, probably due to the dimension space. In our article, we focus on the interest of both positivity and sum-to-one constraints, that allow to estimate the optimum of all possibility of mixtures. Without positivity and sum-to-one constraints, singularities may arise due ``linear dependance'' of endmember spectra (\\citep{Roberts_MappingChaparralin_RSoE1998,Dalton_Exogeniccontrolssulfuric_PaSS2013}). We show that most recent algorithms, including both positivity and sum-to-one constraints, handle this difficulty, thus simplifying tremendously the treatment.  Our aim is to show that additional spectra (flat, slope) minimize the differences between reference laboratory spectra and actually observed spectra, by reducing the effect of grain size and continuum level (photometry, surface roughness, compaction). Some differences between observations and reference endmembers may still be present due to variations in composition, texture,... The selection of the reference spectra database requires always special care and depends on the scientific questions addressed, but at least this database is explicitly known.  Once detection has been performed, time consuming inversions must be applied in order to retrieve quantitative estimate of surface properties (exact mineralogy, abundances, grain size, porosity, ...) as proposed by different techniques such MGM (\\citep{Sunshine_MGM_JGR1993,Kanner_LimitMGM_Icarus2007}) or radiative transfer inversion (\\citep{Doute_PSPC_PSS2007,Poulet_MaficOMEGA1theorie_Icarus2009}). Also the particular mineral type, for instance Al/Fe or Mg/Fe phyllosilicate, must be addressed by additional work using specific spectral library.  We propose here to extend the linear supervised detection technique using a new set of constraint (positivity and sum to unity) and a new set of additional spectra in order to improve the ``mixing coefficient'' estimation. We also introduce a new renormalization with the goal to fit the signal, even in the case of very low signal to noise ratio.  We propose to call LinMin our strategy of using: \\begin{enumerate} \\item linear unmixing technique under constraints (positivity, sum to unity/sum lower than unity) \\item additional spectra made of flat and slope (also possibly cosine functions) \\item renormalization by the variance/covariance noise matrix \\item estimation of the error on the ``mixing coefficient \\end{enumerate} In order to validate the LinMin approach, we test the detection limits on synthetic data simulating a regolith mixture, the grain size effect and the Martian aerosols effect. We also validate the method on actual hyperspectral dataset of OMEGA, CRISM and M$^{3}$. Quantitative estimate of the detection limits are computed. ",
        "conclusions": "We propose to incorporate additional spectra (constant, slope, and cosine functions) to build a new supervised algorithm, called LinMin, based on linear unmixing under positivity and sum-to-one constraints in the goal to detect minerals at the planetary surfaces. The main novelty of this approach is to treat ``linearly dependent spectra\" (such constant, slope and cosine functions) that create degeneracies in usual unconstrained unmixing algorithms, because under our constraints, they are not linearly dependent anymore. Usually, the reference spectra database contains linearly independent spectra (simply because they are different) so the degeneracies are very small. Nevertheless, if the reference spectra database contains significantly linearly dependent spectra, some degeneracies may occur. Adding positivity constraint significantly reduces the degeneracies. In the special case of linear subpixel mixing, the user may be interested by the most precise aerial surface proportion. In this case, once the detection step has been done, a second pass of the LinMin algorithm with the detected spectrum only can enhance the surface proportion estimation.  Some hyperspectral images may be subject to ``spike'' noise, with non-gaussian statistics. As stated in eq. \\ref{eq:LinearMixing}, the linear unmixing methods usually assume gaussian noise, so the results may be corrupted in the worst cases. One solution may be to despike the data using average techniques (\\citep{2008LPI....39.2528P,Carter_Automatedprocessingplanetary_PaSS2013}) before the treatment by LinMin.  We validated the usefulness of LinMin to estimate mixing coefficients in the case of linear mixture on synthetic examples. We also tested numerically the most important non-linear effects on the detection limits : aerosols content and grain size change. Both cases are well treated with our modeling. Some minerals with large scale feature, such olivines and pyroxenes, cannot be detected in the case of additional cosine function because their spectral signature can be also fitted with a mixture of cosine. Nevertheless, the use of additional spectra made of constant and slope components improve the detection of all minerals, including olivines and pyroxenes in comparison with no additional spectra Our approach permits to save the continuum fitting step since it is incorporated in the linear unmixing. We also showed that the knowledge of the noise covariance matrix, that can be estimated from dark current or using other techniques is important to assess the detection limits, and in particular the error on mixing coefficients.  We also tested LinMin on three real cases of hyperspectral images from OMEGA, CRISM and M$^{3}$ instruments. All three cases show detections in agreement with previous analysis, validating the LinMin with approach. The main difficulty of this approach, that is also present in band ratio and other detection method, is to optimize the threshold for detection. Nevertheless, this problem is partly tackled by the estimation of the error on the mixing coefficients. A combination of threshold on RMS residues, error on mixing coefficients, and maximum mixing coefficients seems to be the best compromise to ensure automatic detection. This strategy has to be fully validated on large dataset.  IPLS is shown to be the best numerical algorithm to solve the LinMin problem. Its fast GPU implementation is particularly relevant for the treatment of large dataset of hyperspectral images. In the future, this methodology should be applied in various planetary cases in order to study the surface geology, especially in more challenging detection situation such complex mafics and anorthosites assemblage on the Moon (\\citep{Ohtake_globaldistributionpure_N2009}), or mixture of hydrous sulfates, hydrated acid and water ice on Europa (\\citep{McCord_HydratedmineralsEuropa_I2010,Dalton_Exogeniccontrolssulfuric_PaSS2013}). Also a significant improvement of the mineral detection may be addressed by using spectral database adapted to the context.   \\subsubsection*{Acknowledgement}  We acknowledge support from the ``Institut National des Sciences de l'Univers'' (INSU), the ``Centre National de la Recherche Scientifique\" (CNRS) and ``Centre National d'Etude Spatiale\" (CNES) and through the \"Programme National de Plan{\\'e}tologie\"."
    },
    "1402/1402.4766_arXiv.txt": {
        "abstract": "Millisecond duration bright radio pulses at 1.4-GHz with high dispersion measures (DM) were reported by Lorimer et al., Keane et al., and Thornton et al.  Their all-sky rate is $\\approx 10^4$/day above $\\sim$1 Jy.  Related events are ``Perytons'' -- similar pulsed, dispersed sources, but most certainly local.  Suggested models of fast radio bursts (FRBs) can originate in the Earth's atmosphere, in stellar coronae, in other galaxies, and even at cosmological distances.  Using physically motivated assumptions combined with observed properties, we explore these models.  In our analysis, we focus on the Lorimer event: a 30 Jy, 5-ms duration burst with DM$=$ 375 cm$^{-3}$ pc, exhibiting a steep frequency-dependent pulse width (the {\\it Sparker}). To be complete, we drop the assumption that high DMs are produced by plasma propagation and assume that the source produces pulses with frequency-dependent arrival time (``chirped signals''). Within this framework we explore a scenario in which Perytons, the {\\it Sparker}, and the FRBs are all atmospheric phenomenon occurring at different heights. This model is {\\it ad hoc} in that we cannot explain why Perytons at higher altitudes show greater DMs or exhibit narrower pulses.  Nonetheless, we argue the {\\it Sparker} may be a Peryton.  We end with two remarks. First, the detection of a single FRB by an interferometer with a kilometer (or longer) baseline will prove that FRBs are of extra-terrestrial origin.  Second, we urge astronomers to pursue observations and understanding of Perytons since they form (at least) a formidable foreground for the FRBs. ",
        "introduction": "\\label{sec:Introduction}  The subject of radio transients seems to have finally come of age. The Galactic list starts with GCRT~J1745$-$3009, an erratic source in the meter-wave band \\citep{hlk+05} and is followed by neutron stars which produce strong pulses occasionally, the so-called Rotating Radio Transients (RRATs; \\citealt{mll+06}). Radio afterglows appear to routinely follow giant flares from soft-gamma repeaters \\citep{fkb99}.  Recently, an entirely new type of radio source \\citep{zbs+11} was unexpectedly discovered first in the hard X-ray band, Swift\\,J164449.3+573451 \\citep{bkg+11}. To within the exquisite astrometric precision afforded by radio VLBI the radio counterpart coincides with the nucleus of a small star-forming galaxy at $z=0.35$ \\citep{ltc+11}.  The Lorentz factor of this relativistic explosion, $\\sim10$, is smaller by an order of magnitude to those inferred in gamma-ray bursts.  A plausible model for the source is blazar activity initiated by feeding a tidally disrupted star to a nuclear black hole \\citep{bgm+11}.  A recent summary of the rates of extra-galactic radio transients can be found in \\citet{fko+12}.  \\citet{lbm+07}  reported the discovery of an intense (30\\,Jy) and short duration (5-ms) burst in the decimeter band (1.4\\,GHz). This transient was found as a result of the archival analysis of the Parkes Multi-Beam pulsar data obtained towards the  Magellanic Clouds.  The dispersion measure (DM) of this burst, 375\\,cm$^{-3}$\\,pc, considerably exceeded  the sum of the estimated dispersion measure contributed by the interstellar medium of our own Galaxy and that contributed by the Magellanic Clouds. Follow-up observations did not show any repeating burst.  Lorimer and colleagues proposed that most of the dispersion measure arose from electrons in the inter-galactic medium (IGM). Using currently accepted values for the density of the IGM they estimated the redshift to this event -- which, here after, we call the {\\it Sparker} -- to be about 0.12 ($\\sim 500$\\,Mpc).  The {\\it Sparker} would be the first impulsive radio transient event seen from outside the Local Group. If so, this discovery assumes a seminal stature.  Specifically, the sharp pulse will enable astronomers to probe the column density, magnetic field and turbulence of the IGM \\citep{mk13,Cordes2013}.  The discovery appeared timely given that several countries have undertaken massive investments in radio astronomy in the meter and decimetric bands --  the Low Frequency Array (LOFAR; \\citealt{fvd+06}), the  Murchison Wide Field Array (MWA; \\citealt{bwk+07}) and the Australian Square-Kilometre-Array Pathfinder (ASKAP; \\citealt{jtb+08}).  In short, in the parlance so popular with funding agencies, the discovery reported by \\cite{lbm+07} could be a transformational finding.  The discovery of the {\\it Sparker} motivated further archival searches.  Additional transients were found with some features similar to that of the {\\it Sparker} \\citep{bbe+11}. Some of these bursts exhibited a trajectory in a plane of  arrival-time ($t$) and frequency ($\\nu$) as follows: $t(\\nu)\\propto \\nu^{-n}$ but with $n\\approx 2$. We remind the reader that for a pulse traveling through cold plasma, $n$ is exactly 2 \\citep{rl79}.  Furthermore, some of the bursts showed ``lumpy'' emission (that is, the broad band spectrum could not be described by a simple power law).  Most troubling was that these events were detected in many beams.  These sources were dubbed {\\it Perytons} by the the discoverers.  Burke-Spolaor and colleagues argue that Perytons are atmospheric phenomenon and explain the detection in all (most) beams to pickup by distant side lobes. \\cite{kbb+12} found additional Perytons in a second re-analysis and noted that a cluster of Perytons were separated by about 22\\,s and suggested that Perytons are artificial signals.  Regardless, it is now accepted that Perytons are terrestrial in origin. The DM of the {\\it Sparker} was noted to be similar to that inferred for Perytons, and as a result some doubt was cast on the extra-galactic nature of the {\\it Sparker}.  Another archival search of the Parkes Galactic Plane Survey data found a transient event in a single beam and with a dispersion measure of 746\\,cm$^{-3}$\\,pc \\citep{ksk+12}.  Earlier this year, \\cite{tsb+13} reported the finding of four short duration bursts. One of these events showed a frequency dependent arrival time with $n=2$ to within the precision offered by the measurements.  These bursts with peak fluxes of about 1\\,Jy also showed dispersion measures (ranging from 553\\,cm$^{-3}$\\,pc to 1104\\,cm$^{-3}$\\,pc) in considerable excess of that expected from the Galactic interstellar medium.  Unlike the Perytons these four events were found in only one beam.  This archival analysis drew data from the ``High Time Resolution Universe'' (HTRU) survey which in turn used a digital filter bank \\citep{kjv+10}, whereas the older Parkes data were obtained with an analog filter bank.  Thornton and colleagues, like \\cite{lbm+07} before, argue that the excess dispersion arose primarily in the IGM and infer redshifts ranging from $z=0.45$ to $0.96$. These authors quote an all-sky rate of $\\dot{\\mathcal{N}}\\approx 10^4\\,{\\rm events\\ day}^{-1}$.  This is a remarkably high rate for an extra-galactic population (assuming no repetitions).   Curiously, the brightest burst in \\citet{tsb+13} exhibits an asymmetric pulse shape, with a rise time smaller than the decay time.  Furthermore, for this event and the {\\it Sparker}, the observed pulse width is frequency dependent with the pulse width, $\\Delta\\tau$ $\\propto \\nu^{-m}$ and $m\\approx 4$.  Such a characteristic pulse frequency-dependent broadening is also seen in pulsars with large DMs and attributed to multi-path scattering as the radio pulse traverses through inhomogeneous structures in the interstellar medium. In contrast, the Perytons show symmetrical pulse profiles.  To summarize, analysis of the Parkes Multi-Beam data with two different pulsar backends (one analog and the other digital) taken during the course of pulsar searches in the 1.4\\,GHz band at the Parkes Observatory have shown three  types of impulsive radio bursts: Perytons \\citep{bbe+11}, the {\\it Sparker} \\citep{lbm+07} and Fast Radio Bursts \\citep[FRBs,][]{ksk+12,tsb+13}.  There is agreement that Perytons are of terrestrial origin. In contrast, the {\\it Sparker} and FRBs have been argued to arise from extra-galactic sources. ",
        "conclusions": "\\label{sec:Conclusions}   From analysis of archival pulsar data obtained at the Parkes Observatory astronomers have reported radio pulses with millisecond duration and with a frequency dependent arrival time which if interpreted as due to propagation would require dispersion measures considerably exceeding that expected from the Galactic interstellar medium \\citep{lbm+07,tsb+13}. The short durations of these events require a high brightness temperature, even if their origin is Galactic, let alone extragalactic. The all-sky rate of FRBs is an astounding $10^4\\,$ per day.  In this paper we have explored a wide range of scenarios capable of explaining the properties and suggested progenitors of the {\\it Sparker} \\citep[which has the lowest inferred DM;][]{lbm+07} and Fast Radio Bursts (FRBs; \\citealt{tsb+13}). Complicating this discussion is the presence of ``Perytons'' which share the properties of FRBs but are conclusively identified as arising locally (terrestrial origin). The inferred DMs of the Perytons are strongly clustered in the range 300--400\\,cm$^{-3}$\\,pc.   We started our investigation of these sources by accepting that the large inferred DM for the {\\it Sparker} is indeed a result of a signal propagating through a cold plasma.  We arrived at the following conclusions. \\begin{enumerate} \\item Based on available archival imaging, the nebula that produces the large DM for the {\\it Sparker} can be no closer than 300\\,kpc. The minimum distance for the four FRBs (with their larger inferred DMs) would be higher. This conclusion led us to investigate extra-galactic models for the sources.  \\item We consider a host of plausible extra-galactic progenitors including supernovae, blitzars, short hard bursts, white dwarf mergers and soft gamma-ray repeaters. The models either are physically inconsistent (lack a suitable clean and relativistic environment to produce high brightness temperature bursts or suffer from free-free absorption in the general vicinity of the progenitor) or are unable to account for the high all-sky FRB rate ($10^4\\,$day$^{-1}$).  \\item Of all the possible progenitors, giant flares from young magnetars present the most attractive physical model. This model has the advantage of naturally explaining why some FRBs show frequency-dependent pulse widths. The model can also account for the rates provided that an efficiency of $10^{-5}$ can be achieved in converting the energy release in giant flares into radio emission. \\end{enumerate}  We believe that we have explored all reasonable {\\it stellar} models for FRBs. Thus, should it turn out that FRBs are not of stellar origin then non-stellar models (e.g., quasars - E.~S.~Phinney, pers. comm.; cosmic superconducting strings - \\citealt{Vachaspati08}) have to be considered.  Consistent with our agnostic exploration of the FRB phenomenon we drop the requirement that the {\\it Sparker}'s large DM was produced by propagation through a cold plasma. In this framework the source produce a ``chirped'' signal (frequency dependent arrival time). Chirped signals are used by the military (radar), communications (spread spectrum) and also arise from natural phenomenon (e.g., bursts from the Sun, atmospheric events). We propose an empirical model unifying Perytons with FRBs with the Perytons being in the near-field of the Parkes telescope (where the Fresnel approximation holds) and FRBs being in the far-field (where the traditional Fourier optics assumed by radio astronomers holds).  The inferred DM for the {\\it Sparker} is similar to the mode of the Peryton distribution (Figure~\\ref{fig:PerytonHistogram}).  Next, it is not obvious to us (from the signal level in different beams) that the {\\it Sparker} has to be a source at a very large distance (\\S\\ref{sec:BetterLocalization}).  Economy of hypotheses lead us to suggest that the {\\it Sparker}  itself is a Peryton that occurred at a height of about 20\\,km (the Fresnel scale for the Parkes 64-m telescope at a wavelength of 21\\,cm). In order to explain FRBs as Perytons we require that the chirp rate of Perytons must scale proportionally with their distance (height). We offer no explanation for this requirement.  Perytons form a formidable foreground for FRBs. As such further progress will require astronomers to understand the distribution of and distances to Perytons.  Perytons show clearly that Nature can produce chirped signals in the 21-cm band and so a thorough understanding of the Perytons will only help astronomers distinguish local sources from cosmic sources. Since Perytons are local sources with as yet unknown distances some care is needed prior to comparing the rates of Perytons from different telescopes (with differing Fresnel scales).    In summary, there is no compelling evidence to support an extra-terrestrial origin for FRBs. A plausible argument can be made to relate giant flares from SGRs to FRBs.  In this picture the typical redshift of an FRB is $z\\approx 0.5$.  An interferometric localization of FRBs will immediately rule out a local origin. The same data will either show a host galaxy (which would then revive stellar models or quasar models) or no host galaxy (which will favor truly exotic origins). A modest investment in several clusters of simple dipoles tuned to the 1--2\\,GHz band and separated  moderately (tens to hundreds of kilometers) would be a worthwhile investment (if only to explore strong decimetric pulses  not only from Galactic giant flares but from the gamut of Galactic sources).  Despite the current murky situation it is it tempting to think of bountiful diagnostics that can be provided by millisecond bursts of extra-galactic origin.  In \\cite{Zheng14} we review a couple of these diagnostics. In particular we draw the reader's attention to a unique way by which astronomers can search for solar-mass intergalactic MACHOs through FRBs.  We conclude by noting that in the title of the paper: ``Giant Sparks at Cosmological Distances?\", the adjective giant refers to the nominal length scale of the emitting region (300\\,km; \\S\\ref{sec:ExtraGalacticOrigin}), and the word spark has the same meaning as in pulsar phenomenology.  We point out that the traditional outcome of papers that pose a question in their title is generally in the negative. Nonetheless, one could take some comfort from the history of gamma-ray bursts (GRBs). This was an exotic phenomenon even for astronomers.  The history of GRBs started with searches for possible terrestrial (artificial) signals.  Since their discovery in 1967, the diversity of observed phenomena has grown tremendously. Bursts of gamma-rays are now seen from atmospheric events \\citep{fbm+94}, from the Sun \\citep[Third Orbiting Solar Observatory,][]{krau72}, from compact stellar sources in our Galaxy \\citep{mgi+79,cdp+80,kck+08}, and from cosmological distances \\citep{mdk+97} and that too from at least two distinct populations \\citep{kmf+93}.  So, at early times, what one could have  considered to be a single phenomenon literally spans terrestrial to cosmological scales. It may well be that astronomers are on a similar adventure in the radio band."
    },
    "1402/1402.1005_arXiv.txt": {
        "abstract": "{ We present the Green Bank 100\\,m Telescope (GBT) mapping observations of CS 1-0, and the Submillimeter Array (SMA) 157-pointings mosaic of the 0.86 mm dust continuum emission as well as several warm and dense gas tracers, in the central $\\sim$20 pc area in Galactic Center. The unprecedentedly large field-of-view and the high angular resolution of our SMA dust image allow the identification of abundant 0.1-0.2 pc scale dense gas clumps. We found that in the Galactic Center, the Class I methanol masers are excellently correlated with the dense gas clumps. However, on the ~0.1 pc scale, these dense gas clumps still have a extremely large linewidth (FWHM$\\sim$10-20 km\\,s$^{-1}$). Simple calculations suggest that the identified clumps can be possibly the pressurized gas reservoir feeding the formation of 1-10 solar-mass stars. These gas clumps may be the most promising candidates for ALMA to resolve the high-mass star-formation in the Galactic center. } ",
        "introduction": "\\label{sec:back} The universality of the star formation laws has been widely investigated in extragalactic studies ([1], and references therein). When taking a closer look, these star formation laws may represent merely the weighted media for the large number of observed molecular clouds. Ultimately, the high precision observations may allow us to parametrize the details about how the star formation rate and efficiency are deviated from the laws, because of the different physical conditions in the natal molecular clouds. Preliminary hints may first be provided by observing the environments with extreme physical conditions.  The $\\sim$200 pc Galactic Central Molecular Zone (CMZ) provides the unique laboratory to resolve the star-forming molecular clouds around the supermassive black hole with the sub-parsec scale resolutions, thanks to the proximity. The high molecular gas temperature and the largely non-virialized gas motions in these clouds present the most extreme environment for star formation ([2], and references therein). The previous observations on molecular cloud M0.25+0.01 in the CMZ have reported the $\\sim$2 orders of magnitude smaller star-forming efficiency than the molecular clouds in the rest of the Galactic disk [3, 4, 5, 6]. In this work, we present the multi-scale mapping observations towards the Sgr\\,A molecular clouds, which are the star-forming regions nearest to the central supermassive black hole Sgr\\,A*. Our aims were to resolve the overall geometry and kinematics of these molecular clouds, and to examine the formation of the localized gas overdensities in high resolution images. Our large-scale mapping observations have provided the road maps for the higher angular resolution follow-ups.  \\vspace{-0.3cm} \\begin{figure}[h] \\hspace{-0.9cm} \\includegraphics[width=12.5cm]{arc.png} \\vspace{-0.5cm} \\caption{\\small{ The channel maps of the CS 1--0 line (gray scale; [7]). The color bar is in Kelvin units. White contours show the 300\\,Jy\\,beam$^{-1}$\\,km\\,s$^{-1}$ level of the velocity integrated HCN 4--3 intensity map taken by SMA, to indicate the warm CND. The orange dashed-dotted arc is drawn to indicate the newly discovered Southern Arc. The light blue dashed-dotted arc is drawn to indicate the Eastern Arm constituted by the well-known 50 km\\,s$^{-1}$ molecular cloud, the northern part of the 20 km\\,s$^{-1}$ molecular cloud, and the molecular ridge located in between. }} \\label{fig_arc} \\vspace{-0.8cm} \\end{figure} ",
        "conclusions": ""
    },
    "1402/1402.0003_arXiv.txt": {
        "abstract": "Using the UltraVISTA catalogs, we investigate the evolution in the 11.4~Gyr since $z=3$ of the progenitors of local ultra-massive galaxies ($\\log{(M_{\\rm star}/M_{\\odot})}\\approx11.8$; UMGs), providing a complete and consistent picture of how the most massive galaxies at $z=0$ have assembled. By selecting the progenitors with a semi-empirical approach using abundance matching, we infer a growth in stellar mass of 0.56$^{+0.35}_{-0.25}$ dex, 0.45$^{+0.16}_{-0.20}$~dex, and 0.27$^{+0.08}_{-0.12}$ dex from $z=3$, $z=2$, and $z=1$, respectively, to $z=0$. At $z<1$, the progenitors of UMGs constitute a homogeneous population of only quiescent galaxies with old stellar populations. At $z>1$, the contribution from star-forming galaxies progressively increases, with the progenitors at $2<z<3$ being dominated by massive ($M_{\\rm star} \\approx 2 \\times 10^{11}$M$_{\\odot}$), dusty ($A_{\\rm V}\\sim$1--2.2 mag), star-forming (SFR$\\sim$100--400~M$_{\\odot}$ yr$^{-1}$) galaxies with a large range in stellar ages. At $z=2.75$, $\\sim$15\\% of the progenitors are quiescent, with properties typical of post-starburst galaxies with little dust extinction and strong Balmer break, and showing a large scatter in color. Our findings indicate that at least half of the stellar content of local UMGs was assembled at $z>1$, whereas the remaining was assembled via merging from $z\\sim 1$ to the present. Most of the quenching of the star-forming progenitors happened between $z=2.75$ and $z=1.25$, in good agreement with the typical formation redshift and scatter in age of $z=0$ UMGs as derived from their fossil records. The progenitors of local UMGs, including the star-forming ones, never lived on the blue cloud since $z=3$. We propose an alternative path for the formation of local UMGs that refines previously proposed pictures and that is fully consistent with our findings. ",
        "introduction": "One of the most controversial questions regarding the formation and evolution of galaxies is when and how today's most massive galaxies formed. In the standard paradigm of structure formation, dark matter halos build in a hierarchical fashion through the dissipationless mechanism of gravitational instability, while galaxies form inside these structures following the radiative cooling of baryons, with stars forming out of gas that cools within the halos. In this picture, the most massive galaxies are naturally at the endpoint of the hierarchical merging process, and offer an opportunity to constrain models of galaxy formation and evolution, since their predicted galaxy properties are sensitive to various model assumptions, such as gas cooling, star formation, stellar and AGN feedback processes, chemical evolution, and mergers (e.g., \\citealt{baugh06}; \\citealt{delucia06}; \\citealt{henriques12}; and references therein).  In the local universe ($z\\sim0$), the most massive galaxies (i.e., stellar mass $M_{\\rm star}>5\\times10^{11}$~M$_{\\odot}$) constitute a homogeneous population. These local ultra-massive galaxies tend to live in high-density environments (typically rich groups or clusters of galaxies; e.g., \\citealt{blanton09} for a review) and are characterized by early-type morphologies, red rest-frame optical colors, extremely quiescent stellar populations, old mean stellar ages, significant enhancements of the $\\alpha/Fe$ element ratio, and little dust attenuation, and populate the extreme massive end of the very tight red sequence (e.g., \\citealt{bower92}; \\citealt{kauffmann03}; \\citealt{kauffmann04}; \\citealt{gallazzi05}; \\citealt{nelan05}; \\citealt{thomas05}; \\citealt{baldry06}; \\citealt{gallazzi06}; \\citealt{thomas10}). In particular, two decades of archeological studies (i.e., the detailed investigation of the stellar populations in local galaxies to constrain their formations and evolutions with cosmic time) have robustly shown that the more massive the galaxy is today, the earlier its star formation must have started and more promptly must have subsided, in an apparently ``anti-hierarchical'' fashion of the star formation. Quantitatively, most of the stars in local ultra-massive galaxies must have formed in the first $\\sim3$~Gyr of cosmic history (i.e., $z>2$) through short ($<$1~Gyr, as low as $\\sim0.2$~Gyr for the most massive galaxies), hence intense, bursts of star formation (e.g., \\citealt{renzini06}; \\citealt{vandokkum07}; \\citealt{thomas10}).  Increasingly more sophisticated models of galaxy formation and evolution have shown that the apparent down-sizing in the star formation of ultra-massive galaxies is not in contradiction with the hierarchical paradigm. Recent models do predict `anti-hierarchical' star-formation histories for the most massive galaxies in a $\\Lambda$CDM universe (with most stars formed by $z\\sim3$), whereas the assembly of these galaxies is indeed hierarchical, with $\\sim$80\\% of the final mass typically locked up in a single galaxy only at $z<0.4$ and with a number of effective progenitors as high as $\\sim$5 for the most massive galaxies (e.g., \\citealt{delucia06}; \\citealt{delucia07}).  These predictions are qualitatively in agreement with the nearly constant rest-frame optical luminosity density of red-sequence galaxies since $z\\sim1$ found by the COMBO-17 and the DEEP2 surveys (\\citealt{bell04}; \\citealt{faber07}). This result in fact implies that the stellar mass density in red-sequence galaxies has roughly doubled since $z\\sim1$ driven by a combination of galaxy merging and truncation of star formation of the blue star-forming population. In particular, \\citet{faber07} have used these results in support of a scenario for the formation of massive red-sequence (spheroidal) galaxies involving early mass assembly and star formation (with such progenitors living on the blue cloud), followed by quenching (which migrates the progenitors onto the red sequence) and further dry merging (allowing for additional growth along the red sequence to explain the assembly of the most massive galaxies in the local universe).  However, an independent analysis by \\citet{cimatti06} of the evolution of the rest-frame optical luminosity function of red-sequence galaxies since $z\\sim0.8$ from COMBO-17 and DEEP2 has shown that the massive ($M_{\\rm star}>10^{11}$~M$_{\\odot}$) red-sequence galaxies have actually not grown in mass over the past $\\sim$7 billion years, while the inferred growth by a factor of $\\sim$2 of the stellar mass density of red-sequence galaxies since $z\\sim0.8$ appears due to the build-up of the less-massive red-sequence population. In addition to the anti-hierarchical formation of the stars, this result implies also an anti-hierarchical assembly of the stellar content of the most massive galaxies, in contrast to the predictions from theoretical models.  Recent measurements of the stellar mass function of galaxies out to redshift $z=4$ have intensified the tension between observations and theoretical predictions. Indeed, increasingly wider and deeper near-infrared surveys have provided evidence for the existence of very massive galaxies (i.e., $\\log{(M_{\\rm star}/M_{\\odot})} \\gtrsim 11.5$) when the universe was only 1.5 Gyr old (i.e., $z \\sim 4$), and of their number densities evolving very little in the following $\\sim$3--4~Gyr from $z=4$ to $z\\sim$1--1.5 (e.g., \\citealt{perez08}; \\citealt{marchesini09}; \\citealt{marchesini10}; \\citealt{brammer11}; \\citealt{muzzin13b}; and references therein).  While the stellar mass function of galaxies provides a mean to measure the abundance of a galaxy population at a given time and the overall growth as a function of time of its stellar content, it does not tell us how individual galaxies have evolved and assembled their mass. Ultimately, we would like to be able to connect local ultra-massive galaxies with their progenitors in the early universe to understand how their stellar populations have actually changed over cosmic time and to answer the following outstanding questions. Do local ultra-massive galaxies really form in short and intense bursts of star formation at high redshift? When do they actually stop forming stars? How does their stellar population age? How much mass is assembled over time? Do they really evolve from the blue cloud to the red sequence as suggested by, e.g., \\citet{faber07}?  Needless to say, answering these questions have proved difficult, as it requires the non-trivial task of linking galaxies and their descendants/progenitors through cosmic time, which in turn requires assumptions for how galaxies evolve. However, in recent years, several approaches have been developed to link galaxies across cosmic time. A popular approach easy to implement is to compare galaxy properties at fixed cumulative number density (e.g., \\citealt{wake06}; \\citealt{tojeiro10}; \\citealt{brammer11}; \\citealt{papovich11}; \\citealt{tojeiro12}; \\citealt{patel13}; \\citealt{vandokkum13}; \\citealt{leja13}; \\citealt{muzzin13b}). While this approach has the advantage of relying solely on observations, it does not take into account galaxy-galaxy mergers and scatter in mass accretion histories, both affecting the median cumulative number density of a galaxy population (\\citealt{leja13}; \\citealt{lin13}; \\citealt{behroozi13a}). More advanced comparisons have adopted either semi-analytical or semi-empirical galaxy-halo connections to deduce galaxy evolution from simulated dark matter merger histories (e.g., \\citealt{conroy09}; \\citealt{leitner12}; \\citealt{behroozi13a}; \\citealt{moster13}; \\citealt{wang13}; \\citealt{lu12}; \\citealt{leja13}; \\citealt{lin13}; \\citealt{mutch13}; \\citealt{lu14}; \\citealt{behroozi13b}; and references therein).  In this paper, we investigate the evolution with cosmic time of the progenitors of local ultra-massive ($M_{\\rm star} \\approx 6 \\times 10^{11}$ M$_{\\odot}$) galaxies from $z=3$. The progenitors are selected using both the fixed cumulative number density technique and a semi-empirical approach using abundance matching in the $\\Lambda$CDM paradigm. The stellar population properties of the progenitors are then studied as a function of redshift. This analysis provides a direct test of the star-formation histories predicted from the archeology studies of local ultra-massive galaxies and returns a clear and consistent picture of their evolution in the past 11.4 billion years of cosmic history.  Our paper is structured as follows: The adopted dataset and the selection of the sample of progenitors are described in \\S\\ref{sec-dataset} and \\S\\ref{sec-sample}, respectively. The evolution of the properties of the progenitors of local ultra-massive galaxies is presented in \\S\\ref{sec-results}, while the results are summarized and discussed in \\S\\ref{sec-discussion}. We assume $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, and $H_{\\rm 0}=70$~km~s$^{-1}$Mpc$^{-1}$. A \\citet{kroupa01} initial mass function (IMF) is assumed throughout the paper. All magnitudes are in the AB system. ",
        "conclusions": "\\label{sec-discussion}  In this paper we have investigated the evolution since $z=3$ of the properties of the progenitors of ultra-massive galaxies ($\\log{(M_{\\rm star}/M_{\\odot})}\\approx11.8$) at $z=0$ (UMGs). The progenitors have been selected using the fixed cumulative number density approach and a semi-empirical method using abundance matching in the $\\Lambda$CDM paradigm. When the progenitors are selected with the fixed cumulative number density approach, the stellar mass of the progenitors is seen to grow by a factor of $\\sim2.2$ from $z=3$ to $z=0$, implying that about half of the stellar mass in local UMGs was assembled at $z>3$. If the progenitors are selected with the semi-empirical method using abundance matching, the evolution in stellar mass is much more significant ($0.56^{+0.35}_{-0.25}$~dex), with only about a fourth of the stellar mass in local UMGs having been assembled by $z=3$. Since the semi-empirical approach using abundance matching accounts for merging, whereas the fixed cumulative number density method does not, the former is a more desirable and robust approach to constrain the evolution of the median cumulative number density and therefore to select progenitors. The comparison between the assembly history of local brightest cluster galaxies (BCGs) as predicted from semi-analytic models with the inferred evolution in stellar mass since $z=3$ of the progenitors of local UMGs provides yet another manifestation of the previously highlighted tensions between the observed and model-predicted evolutions of the stellar mass function since $z=4$ (e.g., \\citealt{marchesini09}). Specifically, low-mass galaxies in the stellar mass range $10^{9}-10^{11}$~M$_{\\odot}$ form too early in the models and are too passive at later times \\citep{fontanot09}, with the mass growth of simulated BCGs being dominated by accretion since $z\\sim4$ \\citep{delucia07}, in contrast to what is measured in our study.  The progenitors of UMGs have been separated in quiescent and star-forming galaxies using the rest-frame $U-V$ versus $V-J$ diagram ($UVJ$ diagram). At $z=0.35$, the progenitors of UMGs constitute a homogeneous population with similar rest-frame $U-V$ and $V-J$ colors, typical of quiescent galaxies with old stellar populations (age$\\sim$7~Gyr). As the progenitors are followed to higher redshifts, this population becomes increasingly diversified. At $z<1$, the progenitors are all quiescent, while at $z>1$ the contribution from star-forming galaxies progressively increases, with the progenitors' population being dominated by dusty, star-bursting galaxies at $2<z<3$.  The stellar population properties (i.e., stellar age, dust extinction, star-formation rate, specific star-formation rate, rest-frame absolute $V$-band magnitude, and mass-to-light ratio in the rest-frame $V$ band) have been estimated from the modeling of the UV-to-8$\\mu$m spectral energy distributions and 24$\\mu$m data, and their evolutions with redshift have been investigated. At $z=2.75$, most of the progenitors of UMGs are dusty ($A_{\\rm V}\\sim$1-2.2~mag) star-bursting (SFR$\\sim$100-700~M$_{\\odot}$~yr$^{-1}$; sSFR$\\sim$0.8-4~Gyr$^{-1}$) galaxies with a large range in stellar ages, i.e., 0.3-1.6~Gyr. At the same redshift, the quiescent progenitors have properties typical of young (age $\\sim$0.6-1 Gyr) post-starburst galaxies with little dust extinction ($A_{\\rm V}\\sim$0.4 mag) and strong Balmer breaks, showing a large scatter in rest-frame $U-V$ color ($\\sim$0.2 mag). With decreasing redshift, the fraction of quiescent progenitors increases as the star-forming progenitors quench to build the quiescent population. At the same time, the stellar population of the already quiescent progenitors is clearly seen to age with cosmic time, with their properties changing accordingly. Specifically, their stellar ages and mass-to-light ratio increase with decreasing redshift, and their median SEDs evolve from being characterized by a strong Balmer break to being shaped by the 4000\\AA~break in the rest-frame optical. Moreover, the quiescent progenitors are seen to migrate from the bottom-left to the upper-right corner of the $UVJ$ quiescent box and their rest-frame $U-V$ color to redden, as expected from an aging quiescent stellar population.  It is interesting to note that the stellar properties of the star-forming progenitors do not appear to be changing substantially over the redshift interval $1.5<z<3.0$, and that over this same interval the scatter in the $U-V$ color of the quiescent progenitors remains constant and large ($\\sim$0.17~mag) and it starts to decrease only at $z<1.5$. It appears indeed that $z=1.5$ marks the time by which most of the star-forming progenitors have quenched, as indicated by the very few star-forming progenitors at $1.0<z<1.5$. Consistent with this interpretation, the last star-forming progenitors to quench at $z=1.25$ are characterized by lower levels of star formation and older ages on average compared to the star-forming progenitors at $z>1.5$. By $z=1$, all star-forming progenitors have quenched, and the quiescent progenitors are simply left aging thereafter in the following $\\sim$7.7~Gyr to $z=0$, and some additional growing due to, e.g., minor dry mergers.  Whereas the existence at $z=2.75$ of a small fraction ($\\sim$15\\%) of quiescent progenitors implies that the assembly of the very massive end of the local quiescent red-sequence population must have started at $z>3$, our results suggest it had been mostly assembled between $z=3$ and $z=1.5$, with all the growth by in-situ star formation ceased by $z=1$. Most of the quenching of the star-forming progenitors is seen to happen between $z=2.75$ and $z=1.25$, implying a mean age difference of $\\sim$2.6~Gyr among the UMGs. Our results are in remarkably good agreement with the typical formation redshift of quiescent, red-sequence UMGs in the local universe ($z\\sim1.9$) and their rms scatter in age ($\\approx$1.8-2~Gyr) as derived from archeological studies (e.g., \\citealt{gallazzi06}), and provide a direct probe in the early universe of the results and implications of the fossil records in the population of UMGs today.  \\begin{figure} \\epsscale{1.1} \\plotone{fig9.eps} \\caption{Schematic evolutionary tracks plotted in the color-mass diagram showing possible paths for the formation of massive red-sequence galaxies in the local universe. Track C, proposed by \\citet{faber07}, is a ``mixed'' scenario involving both early and late quenching of blue, star-forming galaxies. The red sequence galaxy arise from the blue galaxy when star formation is quenched (upward black arrows from the blue cloud to the red sequence). Once the galaxy arrives on the red sequence, it may evolve more slowly along it through a series of dry mergers (white arrows). The green arrow represents our newly proposed path for the formation of UMGs in the local universe based on our findings. This new scenario involves in-situ early stellar growth during a short and intense dusty burst of star formation at which point the progenitor manifest as a red, massive, heavily dust-extincted star-bursting galaxy. After quenching, the progenitor continues reddening due to aging and potentially growing in mass via dry mergers, reaching its final position on the color-mass diagram. In this new scenario, the progenitor of local UMGs have never lived on the blue star-forming cloud since $z=3$. The population of red galaxies is seen to change with cosmic time, dominated by dusty star-bursting galaxies at $z>2$ and by quiescent galaxies at lower redshift. \\label{fig9}} \\end{figure}  Finally, we have shown that the progenitors of UMGs have never lived on the blue star-forming cloud in the last 11.4~Gyr of cosmic history, i.e., since $z=3$. This is true also for the star-forming progenitors, which have a median rest-frame $U-V$ color at $z=2.75$ similar to the quiescent progenitors, with a significant fraction of them having $U-V$ colors much redder than most of the quiescent progenitors at the same redshift. In light of our findings, the picture in which early-type UMGs at $z\\sim0$ have formed from star-forming galaxies living on the blue cloud that reddened after quenching and further grew through moderate dry merging has to be refined. Such scenario is pictured in Figure~\\ref{fig9} as taken from the ``Evolution Track C'' in \\citet{faber07}. This scenario involves early mass assembly and star formation in galaxies on the blue cloud of star-forming galaxies. When the star-formation activity in the blue galaxies gets quenched, they rapidly move onto the red sequence. Once on the red sequence of quiescent galaxies, any further growth in stellar mass in these progenitors mostly proceeds via dry merging among galaxies on the red sequence. Whereas such path is still viable to explain the formation of less massive (e.g., $10^{11}$~M$_{\\odot}$) local red-sequence galaxies which may have gone through mostly unobscured (hence blue) episodes of early star formation, our study suggest an alternative path for the formation of UMGs at $z\\sim0$. This new scenario, pictured as a green arrow in Figure~\\ref{fig9}, involves early mass assembly and stellar growth in a short and intense dusty burst of star formation during which the progenitors manifest as a red, heavily dust-obscured star-bursting galaxies. After quenching, the progenitors continue reddening due to the aging of their stellar populations, and eventually grow in mass through (mostly minor) dry merging, finally reaching their current positions on the color-mass diagram. This scenario qualitatively agrees with the predictions from the theoretical model of the co-evolution of massive galaxies and super-massive black holes originally proposed in \\citet{granato04} and further developed in \\citet{cai13} and \\citet{lapi14}, which indicates that the star formation in the progenitors of local very massive spheroids proceeded within a heavily dust-enshrouded medium at $z>1.5$ over a timescale $\\lesssim 0.5-1$~Gyr.  Our study provides a complete and consistent picture of how the most massive galaxies in the local universe have formed and assembled since $z=3$. While we see the star-forming progenitors disappearing and turning into the quiescent progenitors, we do not know what physical mechanisms are responsible for halting the star formation. Is the AGN, if present, playing a major role? What about feedback from the starburst itself? For the first time, our analysis provides a sample of star-forming progenitors that must quench by $z=1$. Detailed studies of the AGN content, physical conditions of the gas (e.g., gas fraction and gas-phase metallicity), stellar population properties (e.g., stellar ages and metallicities, SFRs), structural and dynamical properties, and environment of these star-forming progenitors can shed light into the physical processes responsible for the inevitable quenching of their star-formation activity. High $S/N$ optical and near-infrared spectroscopic studies, as well as deep sub-millimeter studies with ALMA will provide the datasets to answer the remaining outstanding questions.  Needless to say, our study needs to be extended to earlier times to investigate the properties of the progenitors at $z>3$, i.e., in the first 2 billion years of cosmic history. Have the progenitors ever lived on the blue star-forming cloud? Which are the progenitors of the quiescent, massive, post-starburst galaxies we see at $2.5<z<3$? Much deeper near-infrared selected catalogs over a large surveyed area are required to push this investigation to higher redshift. The second data release (DR2) of the UltraVISTA survey will provide much deeper near-infrared coverage over $\\sim$70\\% of the full COSMOS field, reaching $\\sim1.2$~mag deeper in the $K_{\\rm S}$ band, and allowing to probe the stellar population 0.5~dex deeper in stellar mass.  The observed assembly since $z=3$ of the most massive galaxies need to be investigated in more detail in the context of theoretical models of galaxy formation and evolution. The new generation of semi-analytic models accounting for a better treatment of satellite galaxies, for the formation of the intra-cluster component, including detailed chemical enrichment models, and coupled with radiative transfer codes or templates to predict the infrared luminosities of model galaxies (e.g., \\citealt{vega05}; \\citealt{lagos12}) will provide valuable tools to further test the physical mechanisms driving the formation and evolution of today's most massive galaxies.  Finally, we note that $\\sim$19\\% of the star-forming progenitors at $z>2$ have extraordinary SEDs with very red $V-J$ colors ($V-J>1.9$), potentially representing a ``new'' population of galaxies in the early universe that effectively disappear by $z\\sim1$. While unlikely to change our overall picture of the formation of UMGs, systematic effects dominate the uncertainties in both the photometric redshifts and SED modeling of individual galaxies among this extreme population. Detailed investigations, including spectroscopy with near-infrared spectrographs and ALMA, and near-infrared imaging with the {\\it Hubble Space Telescope}, will be needed to confirm the existence of this intriguing population and for a more complete understanding of the evolutionary history of UMGs."
    },
    "1402/1402.0832_arXiv.txt": {
        "abstract": "\\noindent  We present Atacama Large Millimeter/submillimeter Array (ALMA) observations in the dust continuum (690~GHz, 0.45~mm) and $^{12}$CO~$J=6-5$ spectral line emission, of the transitional disks surrounding the stars \\SAO\\ and \\SR. These ALMA observations resolve the dust-depleted disk cavities and extended gaseous disks, revealing large-scale asymmetries in the dust emission of both disks. We modeled these disks structures with a ring and an azimuthal gaussian, where the azimuthal gaussian is motivated by the steady-state vortex solution from \\citet{2013Lyra}. Compared to recent observations of \\HD, \\OphIRS, and \\LkHa, these are low-contrast ($\\lesssim2$) asymmetries. Nevertheless, a ring alone is not a good fit, and the addition of a vortex prescription describes these data much better. The asymmetric component encompasses $15\\%$ and $28\\%$ of the total disk emission in \\SAO\\ and \\SR\\, respectively, which corresponds to a lower limit of $2M_{Jup}$ of material within the asymmetry for both disks. Although the contrast in the dust asymmetry is low, we find that the turbulent velocity inside it must be large ($\\sim20\\%$ of the sound speed) in order to drive these azimuthally wide and radially narrow vortex-like structures. We obtain residuals from the ring and vortex fitting that are still significant, tracing non-axisymmetric emission in both disks. We compared these submillimeter observations with recently published H-band scattered light observations. For \\SR\\, the scattered light emission is distributed quite differently from submillimeter continuum emission, while for \\SAO\\, the submillimeter residuals are suggestive of spiral-like structure similar to the near-IR emission. ",
        "introduction": "The structure of a circumstellar disk is expected to dramatically change during the planet formation process. In a few~million years an optically thick massive protoplanetary disk will transform into an optically thin debris disk with little material \\citep[e.g.][]{2007ApJ...662.1067H}. Processes like viscous accretion, photo evaporation, grain growth, and planet formation shape the structure of circumstellar disks \\citep{2001MNRAS.328..485C,2005A&A...434..971D,2008ApJ...678L..59I}, and in particular, the process of planet formation creates gaps, cavities, and asymmetries that can be directly observed \\citep{2005ApJ...619.1114W}. Circumstellar disks with gaps and/or cavities lack warm dust close to the star while still possessing a significant reservoir of cold dust in the outer disk. Such disks are classified as transitional disks, and plenty is known about their inner disk gaps/cavities \\citep[e.g.][]{2008A&A...490L..15D,2009ApJ...699..330S,2010ApJ...717..441E,2012ApJ...747..136I}. However, it is only recently that observations have reached the sensitivity and angular resolution required to study asymmetrical features in the outer disk structure. Particularly, the only known examples of high contrast asymmetries in transitional disks at submillimeter wavelengths are \\HD\\ \\citep{2013Natur.493..191C,2013PASJ...65L..14F}, \\OphIRS\\ \\citep{2013Sci...340.1199V}, and \\LkHa\\ \\citep{2013ApJ...775...30I}. To understand the process of planet formation, long wavelength observations are critical, since they are less affected by optical depth, and hence trace the bulk of the mass surface density.   In this letter, we present ALMA observations at 0.45~mm of two disks surrounding the stars \\SAO\\ and \\SR, which exhibit large-scale asymmetries in their dust continuum emission that were not known before with high significance. \\SAO\\footnote{also HD~135344B.} is an isolated Herbig Ae/Be star (F4Ve, $1.7\\:M_{\\odot}$), at a distance of 142 pc in the Sco-OB2 association \\citep{2011A&A...530A..85M}. The disk surrounding \\SAO\\ is classified as a transitional disk, with a large dust-depleted cavity ($R_{cav}=46$~AU) and a massive outer disk \\citep{2011ApJ...732...42A,2009ApJ...704..496B}. Recently, scattered light observations of the polarized intensity at near-IR wavelengths revealed spiral-like structure in the outer disk of \\SAO\\ \\citep{2012Muto,2013A&A...560A.105G}. \\SR\\ is a young G3 star ($2.5\\:M_{\\odot}$) in the $\\rho$-Ophiuchus star forming region at a distance of 119~pc \\citep{2008A&A...480..785L}. The dust-depleted cavity  \\citep[$R_{cav}=36$~AU,][]{2011ApJ...732...42A} in this transition disk \\citep{2007ApJ...664L.107B,2009ApJ...704..496B} is not completely devoid of material, as rovibrational CO line observations reveal the presence of inner disk gas \\citep{2008ApJ...684.1323P,2013ApJ...770...94B}. Furthermore, recent polarized intensity H-band observations of \\SR\\ by \\citet{2013ApJ...767...10F}  display emission inside the cavity radius, and unlike \\SAO, the radial profile of emission is quite smooth.   \\begin{figure*}[!t] \\begin{center} \\includegraphics[scale=0.16]{Fig1.pdf} \\caption{ALMA observations of dust and gas emission for the disks surrounding \\SAO\\ (top row) and \\SR\\ (bottom row). Ellipses indicate beam sizes listed in Table~1. Left panels: dust continuum emission (color scale) and dust brightness temperature (contours starting at 10~K, spaced by 5~K). Middle panels: $^{12}$CO~$J=6-5$ moment~0 map, contours start and are spaced by $3\\sigma$, where $\\sigma$ is the RMS noise level in the map. Right panels: $^{12}$CO~$J=6-5$ moment~1 map, contours are spaced by $1\\:\\textrm{km~s}^{-1}$.} \\label{amp_vs_uvdist} \\end{center} \\end{figure*}  \\begin{table*}[] \\vspace{-0.5cm} \\caption{Image properties for dust continuum and spectral line observations.} \\label{table_image} \\vspace{-0.3cm} \\begin{center} \\begin{tabular}{l  c c c c c r c c c r} &\t\t\t&\t\t\t&\\multicolumn{4}{c}{Continuum}\t\t\t\t\t\t\t\t\t& \\phantom{a}\t& \\multicolumn{3}{c}{Spectral line}\t\t\t\\\\ \\cmidrule{4-7} \\cmidrule{9-11} Target \t\t& R.A.$^a$\t& Dec$^a$\t& $S_{\\nu}$$^b$\t& RMS \t\t\t& Beam\t\t& Beam\t\t&& RMS$^c$\t \t& Beam\t\t& Beam\t\\\\ &  (J2000)\t\t& (J2000)\t\t& (Jy)\t\t\t& (Jy beam$^{-1}$)& size\t\t& P.A.\t\t&& (Jy beam$^{-1}$)\t& size\t\t& P.A.\t\\\\ \\midrule SAO~206462\t& 15:15:48.448\t& $-$37:09:16.06\t&  4.2\t\t\t&  0.0020\t\t& $0.27'' \\times 0.19''$\t& 53$^{\\circ}$\t&& 0.11\t\t&$0.29'' \\times 0.22''$\t& 60$^{\\circ}$\t\\\\  SR~21\t\t& 16:27:10.281\t& $-$24:19:12.88\t& 3.3\t\t\t\t&  0.0011\t\t& $0.29'' \\times 0.18''$\t& -81$^{\\circ}$\t&& 0.07\t\t&$0.32'' \\times 0.20''$\t& -79$^{\\circ}$\t\\\\ \\bottomrule \\end{tabular} \\end{center} \\vspace{-0.3cm} $^a$ Phase center from the \\citet{2010AJ....139.2440R} catalog, including proper motion (epoch J2000).  \\\\ $^b$ Flux density integrated inside aperture where emission is above $3\\times$RMS noise level. \\vspace{0.1cm} \\\\ $^c$ RMS noise level in a 0.5~km~s$^{-1}$ channel.  \\vspace{0.3cm} \\end{table*}   \\newpage ",
        "conclusions": "We have observed two transitional disks, \\SAO\\ and \\SR, with high sensitivity using ALMA. These transitional disks display large-scale asymmetries in their dust continuum emission, just as two other recently observed with ALMA \\citep{2013Natur.493..191C,2013PASJ...65L..14F,2013Sci...340.1199V}, and an additional disk observed with CARMA \\citep{2013ApJ...775...30I}. It is interesting to note that of the few such objects observed with high sensitivity and high angular resolution, most exhibit high contrast asymmetrical structures. The lack of previous detections is likely due to phase errors introduced by the atmosphere and the lack of sensitivity of previous observations. This may imply that the process (or processes) that drive these asymmetries is taking place in most transitional disks.  In the case of \\SAO\\ and \\SR, we have demonstrated that a ring does not describe these observations, and a vortex-like asymmetry, in addition to a ring of emission, is a better match. Such asymmetrical structures encompass a significant fraction ($\\sim20\\%$) of the total disk emission, although we found that their contrast is much smaller than in \\LkHa, \\HD, and \\OphIRS. Asymmetries in  \\SAO\\ and \\SR\\ have a narrow radial size and a broad azimuthal extent, similar to the \\LkHa\\ and \\OphIRS\\ asymmetrical structures. If these asymmetries are in fact non-transient vortices, the turbulence inside them needs to be substantial in order to produce the azimuthally-broad asymmetrical features.  Despite being a better fit, the vortex prescription does not reproduce every observed feature. Significant residuals remain that are suggestive of spiral-like structure, in the case of \\SAO\\ some residuals roughly coincide with the spiral arms seen in H-band scattered light. To test the presence of large-scale spiral-like features, and to make sure the observed structures trace the mass surface density of the disk down to the midplane, observations at longer wavelengths are required. Future observations with ALMA at 3~mm, and/or with the Very Large Array at 7~mm, will help in this regard, since such observations trace larger dust grains and will have lower optical depth."
    },
    "1402/1402.7079_arXiv.txt": {
        "abstract": "We investigate the statistical dependence of the peak intrinsic colors of Type Ia supernovae (SN Ia) on their expansion velocities at maximum light, measured from the Si II $\\lambda 6355$ spectral feature. We construct a new hierarchical Bayesian regression model, accounting for the random effects of intrinsic scatter, measurement error, and reddening by host galaxy dust, and implement a Gibbs sampler and deviance information criteria to estimate the correlation. The method is applied to the apparent colors from BVRI light curves and Si II velocity data for 79 nearby SNe Ia.  The apparent color distributions of high (HV) and normal velocity (NV) supernovae exhibit significant discrepancies for $B-V$ and $B-R$, but not other colors. Hence, they are likely due to intrinsic color differences originating in the $B$-band, rather than dust reddening. The mean intrinsic $B-V$ and $B-R$ color differences  between HV and NV groups are $0.06 \\pm 0.02$ and $0.09 \\pm 0.02$ mag, respectively.  A linear model finds significant slopes of $-0.021 \\pm 0.006$ and $-0.030 \\pm 0.009$ mag $(10^3 \\text{ km s}^{-1})^{-1}$ for intrinsic $B-V$ and $B-R$ colors versus velocity, respectively. Since the ejecta velocity distribution is skewed towards high velocities, these effects imply non-Gaussian intrinsic color distributions with skewness up to $+0.3$. Accounting for the intrinsic color-velocity correlation results in corrections to $A_V$ extinction estimates as large as $-0.12$ mag for HV SNe Ia and $+0.06$ mag for NV events. Velocity measurements from SN Ia spectra have potential to diminish systematic errors from the confounding of intrinsic colors and dust reddening affecting supernova distances. ",
        "introduction": "Type Ia supernova (SN Ia) light curves have been used as cosmological distance indicators to trace the history of cosmic expansion, detect cosmic acceleration \\citep{riess98,perlmutter99}, and to constrain the equation-of-state parameter $w$ of dark energy  \\citep{garnavich98b,wood-vasey07, astier06, kowalski08, hicken09b, kessler09, freedman09, amanullah10,conley11,sullivan11,rest14,scolnic14b}.   Determining supernova distances with high precision and small systematic error is essential to accurate  constraints on the cosmic  expansion history and the properties of dark energy. However, the confounding of extrinsic host galaxy dust reddening with the intrinsic color variations of SNe Ia  presents a systematic limitation to their precision and accuracy in cosmological applications \\citep{conley07}.  In this paper, we investigate how measurements of the expansion velocity of the supernova atmosphere can be used to learn more about the intrinsic color distribution of SNe Ia and improve inferences of host galaxy dust extinction.  Dust along the line of sight in the host galaxy reddens SN Ia colors and dims their magnitudes, and can lead to systematic errors in distance estimates if not properly accounted for.  The ratio of total to selective dust extinction, $R_V = A_V/(A_B - A_V)$, typically parameterizes the wavelength dependence of dust absorption and scattering, and has an average value of 3.1 for interstellar dust in the Milky Way (MW) Galaxy, although it can vary between 2.1 and 5.8 \\citep{draine03}.   Studies of external galaxies have found similar extinction curves with $R_V \\approx 2.8$ \\citep{finkelman08, finkelman10}.  However, several analyses of both individual SNe Ia and large samples have found anomalously low effective values of $R_V \\lesssim 1.8$ \\citep{branch92, tripp98,trippbranch99,conley07,nobili08,hicken09b,krisciunas07, elias-rosa06, elias-rosa07, wangx08}.    Interestingly, recent analyses of multi-wavelength light curve and color data, including near-infrared (NIR) observations, have indicated that SNe Ia with low reddening are subject to dust extinction $A_V$ with a reddening law closer to $R_V \\approx 3$, whereas highly reddened objects appear extinguished by dust with $R_V \\approx 1.7$ \\citep{folatelli10, mandel11, burns14}.  Systematic uncertainties in the treatment of dust and color of SNe Ia have important implications for cosmological inference.  For example, \\citet{scolnic14} find that attributing the Hubble diagram residual scatter to SN Ia color dispersion rather than luminosity dispersion yields an effective $R_V$ consistent with MW dust, and yields a 4\\% shift in the inferred value of $w$.  Resolving the confusion between the intrinsic color variation in the SN Ia population and extrinsic host galaxy dust extinction is needed for the proper analysis of SN Ia observables.  One promising strategy is to observe SNe Ia at rest-frame NIR wavelengths, where host galaxy dust extinction is diminished, and SNe Ia have been shown to be excellent standard candles using a burgeoning nearby sample \\citep{elias85, meikle00, krisciunas04a, krisciunas04c, wood-vasey08, mandel09, contreras10, folatelli10, stritzinger11, barone-nugent12, weyant13, friedman14}.  At rest-frame optical wavelengths, where the bulk of nearby and high-$z$ observations of SNe Ia have been made, we must find better ways to decompose the observed apparent magnitudes and colors into the components intrinsic to the SN Ia and those extrinsic and due to dust reddening and extinction.   Current SN Ia analyses using the SALT2 method \\citep{guy07} do not distinguish between intrinsic SN Ia variations and dust effects.  One way to identify intrinsic color variations is to find the portion of SN Ia color correlated with a measurable intrinsic property of SNe Ia, unaffected by dust.  Narrow features in SN Ia spectra are intrinsic to the supernovae and can be correlated with photometric observables, and used as additional parameters to try to improve the estimation of SN Ia luminosities.  Even among ``normal'' SNe Ia there exists diversity in spectral features, such as the widths, strengths, and expansion velocities of specific spectral lines, and their evolution with the phase of the SN Ia.   For example, \\citet{benetti05} subdivided normal SNe Ia into two classes, one with high velocity gradients and one with low velocity gradients.  \\citet{branch06} classified optical spectra of SNe Ia into ``core-normal'', ``broad-line'', ``cool'', and ``shallow-silicon.''   \\citet{blondin12} examined the diversity of SNe Ia in the context of these classification schemes using the large spectroscopic dataset of the CfA Supernova Program.  \\citet{chotard11} modeled the components of the  apparent SN Ia spectroscopic variations depending upon Si II and Ca II H\\&K equivalent widths, finding that the remainder is well described  by a CCM \\citep{ccm89} dust reddening law with $R_V = 2.8 \\pm 0.3$, consistent with the MW value.  The use of spectroscopic variations and their correlations with SN Ia luminosity was explored by \\citet{nugent95,bongard06, foley08} and \\citet{hachinger08}.   \\citet{bailey09} investigated ratios of fluxes at different wavelengths using spectra from the Nearby Supernova Factory and found that they could be used to standardize SN Ia magnitudes and estimate distances with a lower scatter ($\\sim 0.12$ mag) than with light curve information alone.   \\citet*{blondin11} examined the utility of flux ratios and other spectroscopic indicators for predicting SN Ia luminosities and distances using a large spectroscopic data set collected by the CfA Supernova Group.  They used a cross-validation procedure to avoid overfitting the SN Ia sample and to robustly gauge the impact on SN Ia distance estimates in the Hubble diagram.  They found that spectral flux ratios led to modest improvements in the Hubble diagram scatter, but at low statistical significance given the sample size, whereas spectral line profile measurements did not improve luminosity estimates and distance predictions beyond the usual optical light curve width and color standardization.  \\citet{wangx09b}  divided SNe Ia into normal velocity (NV) and high velocity (HV) subclasses based upon the photospheric expansion velocity of the explosion ejecta near maximum light, as measured from the prominent Si II $\\lambda$6355 absorption feature.  Measured from the blueshift of the Si II line, the ejecta velocities are conventionally negative (towards the observer), so HV objects  have more negative velocities ($v < -11,800 \\text{ km s}^{-1}$), and larger absolute velocities $|v|$, whereas NV objects have less negative velocities ($v > -11,800 \\text{ km s}^{-1}$), and smaller absolute velocities $|v|$.  Applying this classification scheme to 156 SNe Ia observed by the Lick Observatory Supernova Search \\citep{ganeshalingam10,silverman12}  and the CfA Supernova Program  \\citep{matheson08}, they analyzed the $B$- and $V$-band magnitudes, light curve shapes and colors of each class.  They found that the Hubble diagram scatter could be reduced from 0.178 mag for the full sample to  $0.12-0.13$ mag by treating each class separately.  They estimated the dust extinction law slope parameter $R_V = 1.9$ for the full sample, but $R_V = 2.4$ for the normal velocity sample alone, and $R_V = 1.6$ for the high velocity sample alone.   Finding that high velocity SNe Ia had redder apparent colors overall, they suggested that treating these subsamples separately with different $R_V$s for host galaxy dust would lead to improved SN Ia distance measurements.  \\citet{foleykasen11} re-analyzed the data set from \\citet{wangx09b} and found a reddening law slope $R_V \\approx 2.5$ fit each subsample separately, when the reddest SNe Ia were excluded ($E(B-V) < 0.35$ mag).  They examined the relation between the $B_\\text{max} - V_\\text{max}$ pseudocolors and the velocity classes, finding that the cumulative distribution function  of the apparent $B_\\text{max} - V_\\text{max}$ pseudocolors of high-velocity SNe Ia were offset by $\\sim 0.07$ mag to the red relative to the CDF of normal SNe Ia, and that the apparently bluest SNe Ia in the normal set were $\\sim 0.10$ mag bluer than the bluest of the high-velocity SNe Ia.  They argued that the offset in the distributions of apparent colors between the two classes is due to an intrinsic color difference between normal-velocity and high-velocity SNe Ia, and found that the rms Hubble residual could be reduced to 0.11 mag by excluding high ejecta velocity SNe Ia.   Hence, accounting for this velocity-color relation (VCR) could lead to better SN Ia distance estimates.  \\citet*{foleysanderskirshner11} examined spectra and color measurements of low redshift SNe Ia observed by the CfA Supernova Program \\citep{blondin12}, and found correlations between the maximum light Si II $\\lambda$6355 velocity and pseudo-equivalent width and the scatter about the mean relation between the peak absolute magnitude in $V$, controlling for light curve shape, and the apparent $B_\\text{max}-V_\\text{max}$ color at peak, which they interpret as due to variations of the intrinsic $B_\\text{max}-V_\\text{max}$ colors.   Using 65 SNe Ia, they estimate a slope of this quantity versus velocity of $-0.033$ mag $(10^3 \\text{ km s}^{-1})^{-1}$, finding that lower velocity SNe Ia tended to have bluer intrinsic colors.  In a similar analysis, \\citet{foley12a} studied spectra and colors of high-$z$ SNe Ia from the SDSS-II and SNLS surveys,  finding similar correlations of ejecta velocity and intrinsic colors.  \\citet{blondin12} presented the spectroscopy from the CfA Supernova Program and compared the maximum light Si II velocities with the intrinsic $B-V$ colors inferred from the \\textsc{BayeSN} statistical model for optical and NIR SN Ia light curves \\citep{mandel11}.   This comprehensive statistical model analyzes the apparent light curve data, incorporating uncertainties due to peculiar velocities, measurement error, host galaxy reddening and extinction, and the intrinsic population distribution to infer intrinsic colors, luminosities and distances.  Regressing the estimates of intrinsic color (in units of mag) inferred from the model, with their uncertainties, against the Si II ejecta velocities (measured in units of $\\text{ km s}^{-1}$) at maximum light, they found a linear slope of $-0.013 \\pm 0.005$ mag $(10^3 \\text{ km s}^{-1})^{-1}$ at 2.9$\\sigma$ significance using 73 SNe Ia.  Analyzing the Carnegie Supernova Project data, \\citet{folatelli13} estimated a similar slope between intrinsic $B_\\text{max} - V_\\text{max}$ and Si II velocity of $-0.012 \\pm 0.016$ mag $(10^3 \\text{ km s}^{-1})^{-1}$, but at lower statistical significance.  The lower significance is not surprising as they used only a small ``low-reddening'' sample   of 9 NV and 4 HV SNe Ia.  The physical origin of the heterogeneity of SN Ia ejecta velocities is still somewhat uncertain.  \\citet{foleykasen11} offer a simplistic, heuristic physical explanation for the correlation between peak intrinsic $B-V$ color and Si II velocity.   Higher ejecta velocity is associated with broader Fe-group absorption lines and thus greater line opacity at wavelengths shorter than $\\sim 4300$ \\AA, in the $B$-band.  The $V$-band is less affected by line opacity and thus by higher ejecta velocities.  Hence, higher ejecta velocity could be correlated with redder intrinsic $B-V$ colors.  Such a trend can be seen in the asymmetric, detonating failed deflagration explosion models of \\citet{kasenplewa07}.  Both the Si II ejecta velocity and the $B-V$ color at maximum are functions of viewing angle into the asymmetric explosion, and there is a roughly linear relation between the two \\citep{foleykasen11}.  \\citet{maeda10} found an association between the early phase velocity gradients, which are correlated with ejecta velocities, and late phase nebular velocity shifts, which probe the inner ejecta of SNe Ia.  The trend supports the hypothesis that their spectroscopic diversity is caused by the effects of viewing angle in observations of asymmetric explosions \\citep{krw09}.  \\citet{maeda11} found correlations between the peak $B_\\text{max} - V_\\text{max}$ pseudocolor, controlling for light curve decline rate, and the nebular emission line shift, further supporting a connection between SN Ia color and viewing angle into an asymmetric explosion.  Correlations between early phase colors and nebular line shift were also found by \\citet{cartier11}. However, \\citet{blondin11b} examined the 2D delayed detonation models of \\citet{krw09} and found that a strong correlation between the $B-V$ intrinsic color and ejecta velocity was not a generic feature of these models.    A recent study by \\citet{wangx13} found that HV SNe Ia tended to occur in brighter regions, closer to the host galaxy center, and within brighter and larger host galaxies as compared to NV SNe Ia.  They suggest that these differences indicate two distinct populations of SNe Ia, possibly associated with different progenitor populations.  Accounting for a significant correlation between intrinsic colors of SNe Ia and their ejecta velocities has important consequences for estimating their dust extinction.   Estimates of dust extinction and reddening in the host galaxies of SNe Ia depend on the statistical properties of the intrinsic colors of SNe Ia, as they are determined from the difference between the observed apparent colors and the inferred intrinsic colors.    Significant correlations of ejecta velocity with intrinsic colors mean that measurement of the ejecta velocity from the SN Ia spectrum would provide additional information to estimate the intrinsic colors of individual SNe Ia more precisely.  This in turn would lead to more accurate inferences of the host galaxy dust extinction that should improve luminosity distance estimates.  Since this effect is not incorporated into current schemes for SN Ia light curve analysis, it has the potential to improve distance estimates.  Although \\citet{wangx09b, wangx13} and \\citet{foleykasen11} split the SNe Ia into two categories based on their Si II velocities, there is no definitive boundary separating the two since the velocities form a continuous distribution.  Since the models of \\citet{kasenplewa07} predict a linear relation between intrinsic color and velocity as a function of viewing angle, it makes sense to construct a statistical model for SN Ia colors with potential linear, or even nonlinear, dependence upon ejecta velocity, treated as a continuous parameter.  In this work, we estimate the functional dependence of multiple intrinsic colors on ejecta velocity by hierarchically modeling the conditional distribution of the observed, apparent optical colors given the measured Si II velocities.  Whereas previous studies have focused on the $B-V$ color, in this work, we estimate the effect in multiple optical colors in the $BVRI$ bands simultaneously, while consistently accounting for dust reddening across these wavelengths.  In \\S \\ref{sec:data}, we describe our dataset of apparent colors from $BVRI$ optical light curves and spectroscopic velocity measurements for a sample of 79 nearby SNe Ia consisting of the compilation of \\citet*{foleysanderskirshner11} plus more recent additional supernovae.   Comparing the apparent color distributions between HV and NV groups, we find statistically significant discrepancies in $B-V$ and $B-R$, but not in other colors.  We demonstrate that these velocity-dependent discrepancies are likely to be caused by intrinsic color differences rather than host galaxy dust extinction.  In \\S \\ref{sec:statmodel}, we construct a new hierarchical Bayesian regression model describing the SN Ia apparent colors and ejecta velocity data as a combination of the mean intrinsic colors-velocity relation, random intrinsic scatter, measurement error, and dust reddening.     The marginal likelihood function, the probability model for the distribution of the observed data depending on the parameters of the population, is described in \\S \\ref{sec:marginal_lkhd} and derived in Appendix \\ref{sec:math_lkhd}.   The empirical distribution of ejecta velocities has a long tail towards high velocities.  Significant correlations of intrinsic colors with velocity generically imply a non-Gaussian marginal population distribution of intrinsic colors. In \\S \\ref{sec:nongaussian}, we demonstrate the capability of our model to capture the resulting non-Gaussian population distributions of the intrinsic SN Ia colors.  In \\S \\ref{sec:globalpost}, the global posterior probability of the unknowns, given the observed data, is derived and depicted as a probabilistic graphical model.     Our model can be used with various hypotheses about the functional form of the relations between the intrinsic colors and ejecta velocity (\\S \\ref{sec:nonlinear}).  We employ  the deviance information criterion \\citep[DIC;][]{spiegelhalter02} to gauge whether the more complex hypotheses are justified by their improved representation of the data.  As it has not been used previously in supernova analyses, we provide a brief introduction to the DIC in Appendix \\ref{sec:modelcomp} for the unfamiliar reader.  We implement a new Gibbs sampling algorithm (Markov Chain for Regressing Colors, \\textsc{MCRC}, Appendix \\ref{sec:mcrc}) to compute the posterior distribution of the unknown intrinsic colors and dust extinctions for individual SNe Ia as well as the population hyperparameters, including those governing the mean relation between intrinsic colors and ejecta velocity.  In \\S \\ref{sec:simulations}, we validate our method on simulated data generated from  non-Gaussian distributions with different underlying intrinsic color-velocity trends to show that the true model is recovered accurately.   The simulations also show that DIC effectively discriminates between competing models for the mean intrinsic color-velocity function.  In \\S \\ref{sec:application}, we apply our new statistical method to our observed data and we find significant trends of intrinsic $B-V$ and $B-R$ colors versus Si II velocity. In \\S \\ref{sec:apply_dic}, using the DIC to evaluate the model fits,  we find that the information criteria significantly favor models with simple non-constant trends over the basic model that assumes a constant Gaussian intrinsic color distribution with no trend with ejecta velocity.  Higher order polynomial fits are disfavored.  In \\S \\ref{sec:implied_color_distr}, we compute the non-Gaussian shape of the population distribution of intrinsic colors implied by the fitted model relations and the ejecta velocity distribution, and estimate its skewness.  In \\S \\ref{sec:effect}, we show that our model capturing the intrinsic colors-velocity trend leads to significant velocity-dependent corrections to intrinsic color and dust extinction estimates for individual SNe Ia.  The accuracy and precision of these estimates are best when using both apparent colors and velocity information.  But even with only the apparent color measurements of an individual SN Ia, one can still  obtain better inferences by using the population distribution of intrinsic colors implied by the model capturing these trends.  This distribution accounts for the skewed probability of intrinsically red events, which is underestimated by the basic Gaussian model that ignores these trends. In \\S \\ref{sec:dust_impact}, we demonstrate the significant velocity-dependent corrections to the dust extinction estimates for the full sample.  We conclude in \\S \\ref{sec:conclusion}. ",
        "conclusions": ""
    },
    "1402/1402.2785_arXiv.txt": {
        "abstract": "We investigate the three types of class~B Bianchi cosmologies filled with a tilted perfect fluid undergoing velocity diffusion in a scalar field background. We consider the two most important cases: dust and radiation. A complete numerical integration of the Einstein field equations coupled with the diffusion equations is done to demonstrate how the presence of diffusion can affect the dynamics of cosmological evolution, where the most attention is paid to changes to the late-time behaviour. We show that aside from quantitative effects, diffusion can result in significant qualitative differences. For example, the cosmologies may recollapse if diffusion is sufficiently strong, or evolve towards the de Sitter state otherwise. In constrast to the diffusionless case, radiation isotropizes in presence of diffusion, and the tilt decreases exponentially at later times:~${V\\sim e^{-0.25\\tau}}$; also, we determine the decay rates of energy density, which become slower when the diffusion term is non-zero. ",
        "introduction": "\\label{Sec:Intro}  In recent years a lot of efforts have been done to investigate diffusional effects both in special and general relativity (see e.\\,g.~\\cite{Herrmann2009,Herrmann2010,Haba2009,Risken1996,Calogero2011,Calogero2012,Calogero2013,Shogin2013,Shogin2014a} and references therein). In cosmology, diffusion is not considered to be a fundamental phenomenon, but rather a model to describe the interaction between two different media. In one of the simplest situations these media can be represented by a perfect fluid and a cosmological scalar field \\cite{Calogero2011,Calogero2012,Calogero2013,Shogin2013,Shogin2014a}. A common assumption for cosmology is that the role of fluid particles is played by galaxies in space. \\par It was shown in \\cite{Calogero2011} how the concept of cosmological constant can be generalized to take diffusion into account. In this case, the cosmological fluid is undergoing velocity diffusion on a scalar field background, which may be associated with cosmic dark energy~\\cite{Calogero2013}. The authors investigated a simple FRW~cosmology and demostrated that presence of diffusion can lead to some interesting effects which do not appear in the standard diffusionless models with a cosmological constant. \\par In \\cite{Shogin2013} we followed \\cite{Calogero2011} and considered a plane symmetric~$G_2$ cosmology with diffusion to get a hint on diffusional effects in spatially inhomogeneous and anisotropic cosmologies. We used the orthonormal frame approach \\cite{Elst1997} and derived the diffusional corrections to the equations governing the dynamics of the universe. The solutions of Einstein field equations coupled with the diffusion equations were studied numerically, and the discovered effects were described. \\par A recent paper~\\cite{Shogin2014a} is devoted to investigation of a general Bianchi type~VIII model with diffusion with an emphasis on analytical methods. In particular, we described some essential features of Bianchi cosmologies with diffusion, found the future attractors and determined the late-time asymptotics of such universes. The current paper is a logical continuation of \\cite{Shogin2013} and can be considered as a satellite paper to~\\cite{Shogin2014a}. In contrast to the latter, in the given work we concentrate our attention on numerical methods of investigation to show how diffusion can affect the evolution process of Bianchi models. \\par The systematic study of anisotropic Bianchi cosmologies started with a pioneering work of Ellis and MacCallum \\cite{Ellis1969}, where the authors applied the general orthonormal frame method to cosmology. Since then, a lot of investigations, both analytical and numerical, have been done. In particular, attention has been paid to cosmologies of types III=VI$_{-1}$~\\cite{Hervik2007,Coley2008},~IV~\\cite{Hervik2005}, and~V~\\cite{Collins1979,Hewitt1992,Harnett1996}, belonging to class~B. It has been shown that in the general case without a cosmological constant these three cosmologies demonstrate completely different behaviour. Also, allowing for a tilted fluid can result in some new interesting phenomena. Considering Bianchi models with tilted fluid taking part in diffusional interactions may be the next step to a deeper understanding of the dynamical properties of anisotropic universes. Bianchi type~V universes, containing the open FRW model~\\cite{Wainwright1997}, are of special interest. \\par The paper is organised as follows. In section~\\ref{Sec:DiffusionModel} we briefly describe the diffusion model and explain how the diffusion is coupled to the Einstein field equations. The dimensionless scale-independent variables are introduced in section~\\ref{Sec:DimensionlessVars}. In section~\\ref{Sec:EqsCons} we write down the equations and constraints governing the dynamics of the considered cosmological models and also discuss the initial conditions. Section~\\ref{Sec:Results} is devoted to results of analytical and numerical investigations. Summary is presented in section~\\ref{Sec:Summary}. ",
        "conclusions": "\\label{Sec:Summary} We have used the dynamical systems approach to perform an analysis of the dynamics of fully tilted dust- and radiation-filled Bianchi models of types~III,~IV, and~V in presence of diffusion. In particular, type~V cosmologies appear to be less subject to the diffusional effects than the other two types under consideration. \\par Aside from the typical quantitative differences which we have described in details, diffusion can drastically change the dynamics of Bianchi cosmologies from a qualitative point of view. The cosmic no-hair theorem is not valid for the models with diffusion, and future recollapse of the universe becomes possible if diffusion is strong enough, in contrast to the standard diffusion-free models with a positive cosmological constant. When recollapse does not happen, the models evolve towards the de Sitter state of accelerated expansion. \\par Another key result is the isotropization of radiation in Bianchi models with diffusion. We have shown that at late times radiation becomes asymptotically non-tilted, while this never happens in the diffusionless case; namely, the tilt components decrease exponentially at later times with~${v_1,~v_2,~v_3\\sim e^{-0.25\\tau}}$. At the same time, the tilt dynamics for dust is not significantly affected by diffusion. \\par Also, we have calculated the decay rates of the energy density in models with diffusion and shown that this quantity decreases slower compared to the case when diffusion is absent: \\begin{equation} \\left\\{ \\begin{array}{lll} \\Omega_{rad} \\sim e^{-4\\tau}, & \\Omega_{dust}\\sim e^{-3\\tau} & \\text{without diffusion}; \\\\ \\Omega_{rad} \\sim e^{-3\\tau}, & \\Omega_{dust}\\to (C_K\\cdot \\tau+C_\\Omega)e^{-3\\tau}, & \\text{with diffusion}. \\end{array} \\right. \\end{equation} On the contrary, the future asymptotic behaviour of the scalar potential and the geometrical variables is not significantly changed in presence of diffusion. \\par The diffusion term asymptotically approaches zero both in the past and in the future, but experiences a maximum at some positive timepoint. At later times the diffusion term has been found to experience an exponential decay,~$K\\sim e^{-3\\tau}$. \\par In type~III cosmologies, diffusion has been shown to affect the future asymptotic value of the $\\lambda$-variable. Namely, the stronger the diffusion, the more significantly~$\\lambda_\\infty$ is reduced. \\par At relatively high initial values of diffusion some changes appear in dynamics of the shear variables in cosmologies of types~III and~IV. The found behaviour is observed in the absence of diffusion only at substantial deviations from the Robertson-Walker-approaching initial conditions. At the same time, this effect has no analogue in Bianchi type~V models."
    },
    "1402/1402.7219_arXiv.txt": {
        "abstract": "Glitches are sudden spin-up events that punctuate the steady spin down of pulsars and are thought to be due to the presence of a superfluid component within neutron stars. The precise glitch mechanism and its trigger, however, remain unknown. The size of glitches is a key diagnostic for models of the underlying physics. While the largest glitches have long been taken into account by theoretical models, it has always been assumed that the minimum size lay below the detectability limit of the measurements. In this paper we define general glitch detectability limits and use them on 29 years of daily observations of the Crab pulsar, carried out at Jodrell Bank Observatory. We find that all glitches lie well above the detectability limits and by using an automated method to search for small events we are able to uncover the full glitch size distribution, with no biases. Contrary to the prediction of most models, the distribution presents a rapid decrease of the number of glitches below $\\sim0.05\\,\\mu$Hz. This substantial minimum size indicates that a glitch must involve the motion of at least several billion superfluid vortices and provides an extra observable which can greatly help the identification of the trigger mechanism. Our study also shows that glitches are clearly separated from all the other rotation irregularities. This supports the idea that the origin of glitches is different to that of timing noise, which comprises the unmodelled random fluctuations in the rotation rates of pulsars. ",
        "introduction": "Neutron stars are the highly-magnetised and rapidly-rotating remnants of the collapse of the cores of once more-massive stars. Having masses of approximately $1.4\\,\\rm{M}_\\odot$ and radii of about $12$\\,km, the high densities of neutron stars indicate a structure of a crystalline-like crust and a superfluid interior \\citep{bpp69,NS1}. Their large and steady moments of inertia mean that they have extremely stable rotational frequencies, which slowly decrease as energy is lost through electromagnetic radiation and acceleration of particles in their magnetospheres. However, this regular spin-down is occasionally interrupted by sudden spin-up events, known as glitches \\citep{rm69,elsk11}.  The exact mechanism responsible for glitches is not fully understood but it is thought to involve a sudden transfer of angular momentum from a more rapidly rotating superfluid component to the rest of the star \\citep{ai75}. This component resides in regions of the interior where neutron vortices, which carry the angular momentum of the superfluid, are impeded in moving by pinning on crustal nuclei or on superconducting vortices in the core (or on both). Since a superfluid in such conditions cannot slow down by outwards motion and expulsion of vortices, the superfluid component will retain a higher rotational frequency as the rest of the star slows down. A glitch occurs when vortices are suddenly unpinned and free to move outwards, allowing for a rapid exchange of angular momentum and the observed spin-up of the crust.  Catastrophic unpinning of vortices is expected once the velocity lag between the two components exceeds a maximum threshold, above which the pinning force can no longer sustain the hydrodynamic lift force exerted on the pinned vortices by the ambient superfluid. It has also been shown \\citep{gl09, agh13} that, beyond some critical lag, a two-stream instability might develop and trigger the unpinning. If in such events the lag is completely relaxed (or partially relaxed by a fixed amount) then the interglitch time interval corresponds to the time it takes for the system to reach the critical threshold again, driven by the nearly constant external torque. Models relying on such a build-up and depletion of the superfluid angular momentum reservoir have been successfully used to explain the regular, similar glitches of some young pulsars \\citep{accp93,pizz11, hps12}. However this simple picture cannot account for the wide range of glitch sizes and waiting times between glitches seen in most pulsars.  Glitch sizes in rotational frequency can range over four orders of magnitude in individual pulsars and appear to follow a power-law distribution \\citep{mpw08}. This favours scale-invariant models of the dynamics of individual vortices in the presence of a pinning potential, such as the vortex avalanche model \\citep{wm08} and the coherent noise model \\citep{mw09}. Alternative models involve non-superfluid mechanisms that can act as unpinning triggers before the critical lag is reached, such as crustquakes \\citep{rud69,bp71} or heating episodes \\citep{le96}. The crustquake-induced glitch model has been particularly favoured for the Crab pulsar as it may explain the persistent changes in slow-down rate observed after some of its glitches \\citep{girp77,accp94} and could possibly lead to a power-law distribution of event sizes, similar to earthquakes.  A second type of irregularity is often seen in the rotational behaviour of pulsars, namely timing noise. Thought to be partially caused by torque variations driven by two or more magnetospheric states \\citep{lhk+10}, it manifests as a continuous and erratic wandering of the rotation rate around the predictions of a simple slow-down model. While glitches are rapid and sporadic events in rotation rate, timing noise appears as a slow and continuous process.  Owing to observational limitations such as infrequent and irregular sampling and the presence of timing noise, the detection of glitches is an uncertain process. Moreover, the signature of timing noise in the data can be confused with glitches, so that the lower end of a glitch-size distribution is possibly contaminated by spurious detections. Knowledge of this distribution is essential for any glitch theory. The largest glitches are easily detected and can be used to constrain the minimum superfluid moment of inertia that can act as an angular momentum reservoir \\citep{aghe12,chamel13}. The biases involved and the question of whether there is a minimum glitch size have not been addressed; so far, the smallest possible glitch has been assumed to lie below our detection limits.  In this paper we study the glitch detection capabilities of the current detection methods and define limits depending upon the intrinsic pulsar rotational stability, observing cadence and sensitivity. We apply these definitions to an extensive set of observations of the Crab pulsar and, by using an automated glitch detector, uncover the full glitch size distribution and show that there is a minimum glitch size. ",
        "conclusions": "\\label{fin} We have quantified our current glitch detection capabilities and, after a meticulous search for small glitches, we have shown that in the case of the Crab pulsar all glitches in this dataset have already been detected. The full glitch size distribution exhibits an under-abundance of small glitches and implies a lower cut-off at $\\Delta\\nu\\sim0.05\\,\\mu$Hz. The existence of such a minimum glitch size implies a threshold-dominated process as their trigger, which still needs to be identified.  Besides the occasional glitches, we have detected a continuous presence of timing noise having a well defined maximum amplitude, which can be described by step changes $|\\Delta\\nu|\\leq 0.03\\,\\mu$Hz and $|\\Delta \\dot{\\nu}|\\leq 200\\times 10^{-15}$\\,Hz\\,s$^{-1}$. The distinct properties of this noise component compared to the glitches imply that timing noise cannot be attributed solely to unresolved small glitches produced by the exact same mechanism."
    },
    "1402/1402.2927_arXiv.txt": {
        "abstract": "{Collapsing magnetic traps (CMTs) have been suggested as one possible mechanism responsible for the acceleration of high-energy particles during solar flares. An important question regarding the CMT acceleration mechanism is which particle orbits escape and which are trapped during the time evolution of a CMT. While some models predict the escape of the majority of particle orbits, other more sophisticated CMT models show that, in particular, the highest-energy particles remain trapped at all times. The exact prediction is not straightforward because both the loss cone angle and the particle orbit pitch angle evolve in time in a CMT.} {Our aim is to gain a better understanding of the conditions leading to either particle orbit escape or trapping in CMTs. } { We present a detailed investigation of the time evolution of particle orbit pitch angles in the CMT model of Giuliani and collaborators and compare this with the time evolution of the loss cone angle. The non-relativistic guiding centre approximation is used to calculate the particle orbits. We also use simplified models to corroborate the findings of the particle orbit calculations.} {We find that there is a critical initial pitch angle for each field line of a CMT that divides trapped and escaping particle orbits. This critical initial pitch angle is greater than the initial loss cone angle, but smaller than the asymptotic (final) loss cone angle for that field line. As the final loss cone angle in CMTs is larger than the initial loss cone angle, particle orbits with pitch angles that cross into the loss cone during their time evolution will escape whereas all other particle orbits are trapped. We find that in realistic CMT models, Fermi acceleration will only dominate in the initial phase of the CMT evolution and, in this case, can reduce the pitch angle, but that betatron acceleration will dominate for later stages of the CMT evolution leading to a systematic increase of the pitch angle. Whether a particle escapes or remains trapped depends critically on the relative importance of the two acceleration mechanism, which cannot be decoupled in more sophisticated CMT models. } {} ",
        "introduction": "Collapsing magnetic traps (CMTs) have been suggested as one of the mechanisms that could contribute to particle energisation in solar flares \\citep[e.g.,][]{SomovKosugi}. The basic idea behind CMTs is that charged particles will be trapped on the magnetic field lines below the reconnection region of a flare. The magnetic field will evolve into a lower energy state, resulting in (a) a shortening of field line length and (b) an increase in the overall field strength. Due to the vast difference in length and timescales between the particle motion and the magnetic field evolution, the particle motion can be described to a large degree of accuracy by guiding centre theory. The conservation of a particle's magnetic moment and the bounce invariant \\citep[e.g.,][]{Keith_grady_2012} give rise to the possibility of an increase in the particle's kinetic energy by betatron acceleration and by first order Fermi acceleration \\citep[e.g.,][]{SomovKosugi,Boga_somov_2005}.  Studies of CMTs using models with varying degree of detail, focusing on different aspects of CMT physics have be carried out \\citep[e.g.,][]{SomovKosugi}. In this paper, we investigate one very important aspect of particle motion in CMTs, namely what determines whether particles remain trapped or escape during the evolution of a CMT. Answering this question is important to be able to assess whether the CMT mechanism can contribute to the energetic particle flux, causing hard X-ray emission from the footpoints of flaring magnetic loops or whether it can be expected to contribute more to emission originating from higher up in the corona \\citep[e.g.,][]{Krucker_2008}.  At first sight, this may seem like a trivial question to answer because obviously as soon as a particle orbit moves into the loss cone it will escape from the CMT. However, on second thought, things are not as simple as they seem for two reasons: (a) the loss cone itself changes in time due to the changing magnetic field strength, and (b) the particle pitch angle also evolves due to betatron and Fermi acceleration. Generally, one would expect the magnetic field strength within a relaxing magnetic loop to increase. Assuming that the magnetic field strength is highest at the foot points and that this field strength does not change substantially, the loss cone should generally open up during the evolution of a CMT. How much the loss cone opens depends on the magnetic field model \\citep[see, e.g.,][for a model where the loss cone opens to $90^\\circ$]{Asch_2004}. Furthermore, whereas in investigations based on a number of simplifying assumptions \\citep[e.g.,][]{Somov_2003,somov_2004,Boga_somov_2005} some predictions can be made about, for example, the relative importance of betatron and Fermi acceleration and the resulting consequences. However, similar predictions are much harder to make for more detailed magnetic field models \\citep[e.g.,][]{Giul2005, Kar_Barta_2006,Mino_2010,Keith_grady_2012} because the relative importance of betatron vs. Fermi acceleration: (i) will be different for particles with different initial conditions; (ii) may change with time during the evolution of the CMT; and (iii) will depend on the details of the magnetic field model used. Furthermore, it is well-known that in time-dependent and curved magnetic fields the two mechanisms are closely linked \\citep[see, e.g.,][]{northrop}.  The aim of this paper is to present a detailed investigation of particle escape and trapping for the CMT model of \\citet{Giul2005}. For the purpose of this paper, we ignore pitch angle scattering by Coulomb collisions, wave-particle interaction, or turbulence. Pitch angle scattering will, of course, change the results, but we regard the present investigation as a benchmark with which possible future investigations, including the effects of pitch angle scattering, can be compared. In Sect. \\ref{sec:model},  we give a brief outline of the CMT model of \\citet{Giul2005} and summarise some of the results of \\citet{Keith_grady_2012} that are relevant to this paper. In Sect. \\ref{sec:pitch}, we investigate how the loss cone evolves in our CMT model and how the pitch angle of typical particle orbits change with time, and we try to give an explanation of our results on the basis of some simplified models in Sect. \\ref{sec:simplemodel}. We present a summary of our results in Sect. \\ref{sec:conclusions} ",
        "conclusions": "\\label{sec:conclusions}  We have presented a detailed investigation of the conditions that affect the trapping and escape of particle orbits in models of collapsing magnetic traps, starting with the investigation of guiding centre orbit calculations in the kinematic MHD CMT model of \\citet{Giul2005}. Based on the results of our investigations and on observations made previously by \\citet{Giul2005} and \\citet{Keith_grady_2012}, we designed a simplified schematic model for CMT particle acceleration and studied three different implementations of this schematic model, which all qualitatively corroborated the previous findings.  We showed that for each magnetic field line in a collapsing magnetic trap there is a critical initial pitch angle, which divides particle orbits into trapped orbits and escaping orbits. This critical initial pitch angle is greater than the initial loss cone angle for the field line, but smaller than the value of the asymptotic loss cone angle for the field line as $t\\to \\infty$.{ We also investigated whether the critical initial pitch angle depends on the initial energy of the particle, and we found that it does not.}  For orbits with initial pitch angle close to the critical value, Fermi acceleration can dominate in the initial phases, but betatron acceleration will become dominant at later times. In the periods where Fermi acceleration dominates over betatron acceleration, the pitch angle will decrease and when betatron acceleration dominates the pitch angle will increase. Due to the nature of CMTs, both mechanisms will always operate simultaneously \\citep[as already stated by][]{Kar_Barta_2006}, but the efficiency of Fermi acceleration has to decrease on a particular field line during the time evolution of a CMT because the motion of the field line must slow down. On the other hand, the magnetic field strength can still continue to increase due to the pile-up of magnetic flux from above.  In the present paper, we have for simplicity excluded the possibility of either Coulomb collisions, wave-particle interactions, or turbulence on the trapping or escape of particles from CMTs. Obviously, each of these mechanisms may change the results we have found in this  paper. The effect of Coulomb collisions on particle acceleration in simple CMT models has been considered by \\citet{Kova_somov_2003}, \\citet{Kar_kos_2004}, and \\citet{Boga_somov_2009}. \\citet{Mino_2011} have investigated the effects of pitch angle scattering by Coulomb collisions on trapping and escape, but without considering the effect of dynamical friction. \\citet{Kar_Barta_2006} included the effects of Coulomb collisions in their MHD model of a CMT. We also mention the recent paper by \\citet{winter_2011}, although a static field is used in that paper.  Obviously, wave-particle interaction and/or turbulence added to CMT model would also change both the energisation process and the pitch angle evolution of particle orbits in a collapsing magnetic trap. As already mentioned by \\citet{Keith_grady_2012}, a combination of a CMT model with stochastic acceleration mechanisms \\citep[see, e.g.,][]{miller} could provide a link between those acceleration mechanisms and the standard flare scenario along the lines proposed by e.g., \\citet{hamilton:petrosian92}, \\citet{park:petrosian95}, \\citet{petrosian:liu04}, and \\citet{liu:etal08}."
    },
    "1402/1402.1115_arXiv.txt": {
        "abstract": "The magnetic reconnection process is studied in relativistic pair plasmas when the thermal and inertial properties of the magnetohydrodynamical fluid are included. We find that in both Sweet-Parker and Petschek relativistic scenarios there is an increase of the reconnection rate owing to the thermal-inertial effects, both satisfying causality. To characterize the new effects we define a thermal-inertial number which is independent of the relativistic Lundquist number, implying that reconnection  can be achieved even for vanishing resistivity as a result of only thermal-inertial effects. The current model  has fundamental importance for relativistic collisionless reconnection, as it constitutes the simplest way to get reconnection rates faster than those accessible with the sole resistivity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.4730_arXiv.txt": {
        "abstract": "We develop and discuss the properties of a new class of lattice-based avalanche models of solar flares. These models are readily amenable to a relatively unambiguous physical interpretation in terms of slow twisting of a coronal loop. They share similarities with other avalanche models, such as the classical stick--slip self-organized critical model of earthquakes, in that they are driven globally by a fully deterministic energy loading process. The model design leads to a systematic deficit of small scale avalanches. In some portions of model space, mid-size and large avalanching behavior is scale-free, being characterized by event size distributions that have the form of power-laws with index values, which, in some parameter regimes, compare favorably to those inferred from solar EUV and X-ray flare data. For models using conservative or near-conservative redistribution rules, a population of large, quasiperiodic avalanches can also appear. Although without direct counterparts in the observational global statistics of flare energy release, this latter behavior may be relevant to recurrent flaring in individual coronal loops. \\changed{This class of models could provide a basis for the prediction of large solar flares}. ",
        "introduction": "Identifying the physical mechanism responsible for coronal heating is still standing as one of the grand challenges of contemporary solar physics \\citep[see][for different theoretical and observational viewpoints]{Klimchuk:2006ka,Aschwanden:2007kq,DePontieu:2011fk}. Nearly a quarter of a century ago, \\citet{Parker:1983hf} began exploring in earnest the possibility that coronal heating could be due to thermal energy released by small magnetic reconnection events occurring continuously throughout the magnetized corona \\citep[see][and references therein for related physical scenarii]{Parker:1983gr,Parker:1988ia,Priest:2002hx}. His physical picture for nanoflare production is now well-known. Photospheric fluid motions associated with turbulent convection randomly shuffle the magnetic footpoints of coronal loop leading to a braiding of fieldlines throughout the volume of the loop, with which is associated an inexorable buildup of current sheets, as individual fieldlines try to return to dynamical equilibrium under the topological constraints imposed by flux-freezing. As these localized current systems build up, plasma instabilities eventually set in and lead to energy release through magnetic reconnection. Parker dubbed these elementary energy release events ``nanoflares'', went on to argue that they occur continuously throughout the corona, and conjectured that they are responsible for coronal heating.   Intense efforts have been expended to observationally verify or contradict Parker's conjecture. This is no easy task, as it involves measuring the energy output of flares very close to the detection limit of the best current solar imagers in space. On the other hand, decades of flare observations have shown that the probability distribution function [$f(E)$] for flare energy [$E$] takes the form of a power law [$f(E)\\propto E^{-\\alpha}$] spanning some eight orders of magnitude in energy, with $\\alpha$ estimated in the range $1.65$ -- $1.95$ \\citep[\\textit{e.g.}][and references therein]{Aschwanden:2000gc,Aschwanden:2002fl,Aschwanden:2011hw}. This is too low for small flares to dominate energy input into the corona (requiring $\\alpha>2$), unless the PDF steepens abruptly below the current detection limit. This is possible but increasingly unlikely as instruments such as the EUV imagers on the \\textit{Transition Region And Coronal Explorer} (TRACE) and \\textit{Solar Dynamics Observatory} (SDO) have pushed the detection limit right up to the boundary of the nanoflare regime. Another possibility is that physical models used to infer thermal energy release from EUV or soft X-ray fluxes are in error, due to erroneous assumptions regarding the intrinsic geometry of the flaring volume \\citep{McIntosh:2001jt}, or of the physical conditions therein, required to infer thermal energy release from observed radiative fluxes \\citep[see][and references therein]{McIntosh:2000kn,Aschwanden:2002jy,Aschwanden:2013kb}.  Whatever recent observational determination of $f(E)$ may imply for coronal heating, Parker's scenario certainly remains a fine model for solar flares in general. Although not originally emphasized by Parker, his nanoflare scenario contains all ingredients now understood to be required to lead to self-organized criticality \\citep[hereafter SOC; see, \\textit{e.g.}][]{Bak:1987ef,Jensen:1998ww}: an open dissipative system -- here a coronal loop -- slowly driven -- here by random footpoint motions -- and subjected to a self-limiting local threshold instability -- here magnetic reconnection. Despite the slow and gradual loading of energy, such a physical system will release energy in an intermittent manner, in the form of avalanches that can span up to the whole system.  \\citet{Lu:1991hp} have proposed an avalanche-type model for solar flares that has become a reference model in the solar context, although numerous variations have now been developed over the intervening years. The present article describes one such variation that (we believe) is particularly interesting in that i) it is readily amenable to (relatively) unambiguous physical interpretation in terms of twisting of a coronal loop, and ii) it can yields qualitatively distinct flaring behavior as the degree of conservation % is allowed to vary. Section \\ref{sec:newmod} briefly describes the original Lu and Hamilton model \\citep[more specifically, the version put forth by][hereafter the ``LH model'']{Lu:1993cy}, along with the various manners in which we have altered the forcing and redistribution rules. Section \\ref{sec:results} provides modelling results, comparing and contrasting the statistical properties of avalanches in the LH model and our variations thereof. We conclude in Section \\ref{sec:conclusion} with a discussion comparing and contrasting our models with other classes of globally and deterministically driven avalanches models developed in different physical contexts, outlining some possible model improvements, and by summarizing the implications of our modelling results for coronal heating (and more specifically with regards to the flaring behavior of small coronal loops) as well as solar flares predictions. ",
        "conclusions": "}  In this article, we have described and documented the behavior of a novel and hitherto unexplored class of lattice-based avalanche models applicable to solar flares. These models, based on a spatially global and fully deterministic driving mechanism, are amenable to physical interpretation in terms of slow twisting of a small coronal loop. Stochasticity is introduced in the models at the level of the redistribution rules and stability threshold governing avalanching dynamics.  Over significant portions of model space, the simulations produce three distinct populations of avalanches. The first is composed of the smallest avalanches with a plateau-like distribution. Then, a classical scale-free distribution of avalanches takes place in the mid-size range of energies. Finally, quasi-periodic avalanches with a more-or-less well-defined mean size make up the third population. The two latter populations are observed to merge in the context of avalanche models incorporating non-conservative redistribution rules.%  We have argued that the large, quasiperiodic avalanches are the reflection of an energy loading/unloading cycle driven by our adopted global deterministic driving mechanism. Even though an avalanche is always triggered at a single lattice node, the effect of the global driver is that energy has been accumulating at {\\it all} internal lattice nodes prior to avalanche onset. Once the avalanche begins, this energy must be either evacuated at the open boundaries, or dissipated locally within the lattice (in non-conservative models). In conservative models, only the former is possible, which leads inevitably to the buildup of very large avalanches that must connect to the boundaries -- and thus end up spanning the whole system -- and remain active until the excess energy has been drained away from the lattice -- thus producing long duration avalanches. In that respect, it is therefore not surprising that a distinct population of large avalanches (and, consequently, break of finite-size scaling) is most readily observed in models based on conservative redistribution rules.  We have been careful thus far to refrain from declaring our models to be in a true self-organized critical state, using instead the term ``SOC-like'' to characterize their avalanching behavior. Formal demonstration of SOC typically requires the computation of critical exponents, and the demonstration that the latter satisfy mutual numerical relations that place them in a specific universality class. % Although interesting in and of itself, the identification of the universality class to which the present model belongs is of limited interest from the point of view of flare modelling. We stress again that the models discussed in this article do satisfy all {\\it sine qua non} physical conditions believed to be required for SOC \\citep{Jensen:1998ww}: a slowly driven open system subjected to a self-limiting, local, finite-threshold instability. Moreover, in nonconservative models characterized by no quasi-periodic large scale events, the lack of dependence of the power-law indices of avalanches size and duration on the driving rate and lattice size certainly suggests that these models all belong to the same universality class.  Although non-conservative redistribution rules have received, to the best of our knowledge, little attention in the solar flare context, true SOC avalanching behavior is known to materialize in non-conservative system. The most studied such SOC system is the cellular automaton formulation of the stick--slip model for earthquake, introduced by \\citet{Olami:1992do}. In that model, non-conservation is related to the fact that a moving block exerts a force not only on its neighbour blocks via the connecting springs, but also on the two plates bounding (and driving) the system from above and below; hence; when an individual block moves, a fraction of the potential energy stored in the springs is dissipated as heat via friction against the bounding plates, rather than all of it ending up in other springs by the end of the avalanche. Something similar also takes place in the non-conservative substorm models of \\citet{Liu:2006cp} and \\citet{VallieresNollet:2010ed}. There, a fraction of the energy released by an unstable magnetic flux tube within the central (equatorial) plasma sheet is lost via MHD waves and/or charged particles travelling away along magnetic field lines back towards Earth, eventually leading to auroral excitation in polar regions. A particularly interesting feature of the latter model is its ability to simultaneously produce ``internal'' avalanches with a scale-free size distribution, and quasi-periodic boundary discharge avalanches with a well-defined mean-size. This squares well with the statistically distinct properties of auroral emission on the one hand, and ring current injection events on the other, both being in principle related manifestation of energy release events in the central plasma sheet. In the model of \\citet{Liu:2006cp}, this effect materialized in part because the boundaries were set up so as to respond dynamically to incoming avalanches; the results presented herein indicate that dual populations of energy release events, each with distinct statistical characteristics, would still materialize in such globally driven systems, even with more conventional ``open'' boundary conditions, provided conservation is high enough. \\changed{Finally, solar flares also exhibit non-conservative properties. Flares are thought to originate from reconnection processes in coronal loops that feed on the stored magnetic energy. Along with the magnetic reorganization associated with a flare, some energy is systematically lost from the system either through radiation (direct X-ray emissions, or hard X-ray emissions triggered by reconnection-accelerated energetic particles), accelerated particles escape or even conduction along the magnetic field lines. We modeled those effects with a simple non-conservation parameter in the avalanching process. Albeit with a quite different cellular automaton model, our results are consistent with the previous finding of \\citet{Hamon:2002hs}: a quasi-cyclic population of large avalanches is found in both cases when the model is close to be purely conservative.}  Notwithstanding the presence or absence of a population of small and large avalanches, in the realm of mid-size avalanches all models are characterized by a power-law distribution of avalanche energy release spanning many orders of magnitude. In general, the corresponding power-law indices are somewhat smaller than for the LH model, implying flatter PDFs, so that all models have a power-law index smaller than the critical value $\\alpha_E=2$ above which the smaller flares (\\textit{i.e.} avalanches) would become the dominant contributors to coronal heating, as in Parker's nanoflare hypothesis.  An obvious extension of the present model is to introduce the third spatial dimension; as argued in Section \\ref{ssec:phys}, the 2D lattices considered herein can be best regarded as a cross section at some fixed position along a coronal loop. Evidently, allowing avalanches to develop along the loop's length would yield a more realistic model. In particular, it would become possible to map the 3D lattice to a bent coronal loop, and therefore compute fractal dimensions of flares/avalanches \\citep[see, \\textit{e.g.} ][]{Morales:2009dc} in the same manner as carried out from observations \\citep[see, e.g.][]{Aschwanden:2013kb}. This would open a new comparative bridge between modelling and observation.  The broad exploration of parameter space made possible with the computationally much less demanding 2D lattices has allowed us to pinpoint the most important parameter regime, which is found for highly dissipative redistribution. Then, the way in which stochasticity is introduced in the model (stability threshold, extraction, redistribution rules) plays a comparatively lesser role in determining the size distribution of avalanches. In other words, the models are robust with respect to the details of stochastic effects.   The present model finally provides an interesting basis for a practical use of avalanches models in the context of solar flares \\citep{Strugarek:2014we}. The deterministic driver minimizes the level of stochasticity embedded in the model, which results in very good predictive capabilities for the larger events. Interestingly, our model could provide an efficient and cheap alternative method for the prediction of large solar flares in the context of space weather.  \\begin{acks} The authors acknowledge stimulating discussions during the ISSI Workshops on Turbulence and Self-Organized Criticality (2012 -- 2013) held in Bern (Switzerland); and during the ``Festival de th\\'eorie'' (2013) held in Aix-en-Provence (France). We acknowledge support from Canada's Natural Sciences and Engineering Research Council. \\end{acks}"
    },
    "1402/1402.6359_arXiv.txt": {
        "abstract": "The Columbia - U. de Chile CO Survey of the Southern Milky Way is used for separating the CO(1-0) emission of the fourth Galactic quadrant within the solar circle into its dominant components, giant molecular clouds (GMCs). After the subtraction of an axisymmetric model of the CO \\back emission in the inner Southern Galaxy, 92 GMCs are identified, and for 87 of them the two-fold distance ambiguity is solved. Their total molecular mass is \\MHH $=$ 1.14 $\\pm$ 0.05 $\\times$ 10$^{8}$ M$\\solar$ accounting for around 40\\% of the molecular mass estimated from an axisymmetric analysis of the \\HH volume density in the Galactic disk \\citep{bron88} \\MHH$_{disk}$ $=$ 3.03 $\\times$ 10$^{8}$ M$\\solar$. The large scale spiral structure in the Southern Galaxy, within the solar circle, is traced by the GMCs in our catalog; 3 spiral arm segments: the \\emph{Centaurus}, \\emph{Norma}, and \\emph{3-kpc expanding} arm are analyzed. After fitting a logarithmic spiral arm model to the arms, tangent directions at 310\\deg, 330\\deg, and 338\\deg, respectively, are found, consistent with previous values from the literature. A complete CS(2-1) survey toward IRAS point-like sources with FIR colors characteristic of UC HII regions is used to estimate the massive star formation rate per unit \\HH mass (MSFR), and the massive star formation efficiency ($\\epsilon$) for GMCs. The average MSFR for GMCs is 0.41 $\\pm$ 0.06 L$\\solar$/M$\\solar$, and for the most massive clouds in the \\emph{Norma} arm it is 0.58 $\\pm$ 0.09 L$\\solar$/M$\\solar$. Massive star formation efficiencies of GMCs are on average 3\\% of their available molecular mass.\\\\ ",
        "introduction": "Within the Galactic disk, at 100 parsecs scales, the molecular gas is contained mostly in the form of giant molecular clouds (GMCs) \\citep{dame86,bron88,williams97}. The large scale clumpy structure of the CO emission in longitude-velocity diagrams is clear evidence of the organization of the gas into these large objects \\citep{bron88b}. The principal characteristic of GMCs is the role they play as tracers of the large scale structure in the Galaxy and as birthplaces of most of the massive stars in the Galactic disk. \\\\  As for the origin of GMCs, it is accepted that, in the disk of the Milky Way, they are formed as the molecular gas enters in the spiral wave pattern of the gravitational potential energy of the Galactic disk \\citep{elmegreen94}. Tidal shear forces among GMCs are weaker within the spiral arms than for the interarm regions making the gravitational collapse more feasible \\citep{luna06}. Other mechanisms, such as the growth by collisions are, apparently, too inefficient to reproduce the observed power law mass distribution of GMCs \\citep{elmegreen93b,tan2012}.\\\\  GMCs are excellent tracers of large scale Galactic structure. The best example is found in the \\emph{Carina} spiral arm, traced over 20 kpc in the outer Galaxy by more than 40 GMCs, between $l =$ 270\\deg\\textrm{ }and $l =$ 330\\deg\\textrm{ } in Galactic longitude \\citep{grabelsky87,grabelsky88}. \\citet{dame86} reconstructed a 3 spiral arm model for the first Galactic quadrant based on their catalog defined in CO and suggested that the containment of the largest molecular clouds in the arms demonstrates that CO emission is enhanced in the arms not only because the clouds are hotter as suggested by \\citet{sanders85}, but mainly because they are larger and contain more mass.\\\\  GMCs are the known places of massive star formation. Massive stars ($M > 8$ M$\\solar$) originate mainly inside dense, compact clumps (UC HII regions)  in giant molecular clouds \\citep{evans99,maclow03,krum05b,tan05,luna06,mckee07,zinnecker07,schuller2010}. At large scales, \\citet{bron00} showed that the radial distribution of massive star forming regions follows, on average, that of the molecular gas in the inner Galactic disk for all galactocentric radii and that the massive star formation is highest at the peak of the southern ``molecular ring'' (understood as an azimuthally averaged molecular gas distribution of the Galactic disk), between 0.5 $\\leq$ R/R$\\solar \\leq $ 0.6 in galactocentric radius. For the fourth Galactic quadrant (\\IV), \\citet{luna06} showed that massive star formation occurs in regions with high molecular gas density, roughly coincident with the line of sight tangent to spiral arms. At smaller scales, giant molecular clouds harbor most of the massive star formation in the Galaxy \\citep{luna06,maclow03,mckee07,zinnecker07}. An example of this is the extremely high velocity molecular outflow, signature of massive star formation \\citep[and references therein]{krum05}, in the G331 region \\citep{bron08,merello2013}.\\\\  While extensive work has been carried out to identify and find the physical characteristics and spatial distributions of GMCs in the first Galactic quadrant (\\I) within the solar circle \\citep{dame86,scoville87,solomon87}, no equivalent catalog of GMCs in the \\IV within the solar circle has been published yet. For the \\I, different methods have been used to define the GMCs: topologically closed surfaces in the three dimensional LBV (longitude, latitude, velocity) CO data phase space  \\citep{scoville87,solomon87}; \\emph{``clipping''} the CO data below a certain temperature threshold \\citep{myers86}; and subtraction from the observed CO emission of a synthetic dataset generated from an axisymmetric model of the Galactic molecular gas distribution (background subtraction; \\citep{dame86}), to ``extract'' the GMCs from the complex CO \\back emission in which they are immersed. In this work, we use  the CO dataset from the Columbia - U. de Chile Southern CO Survey of the Milky Way \\citep{bron88b} (hereinafter referred to as CO Survey) to define GMCs in the 3-dimensional phase space through the subtraction, from the CO data, of an axisymmetric model (ASM; \\citep{bron88}) of the complex CO \\back emission following \\citet{dame86}. \\\\  We estimate the massive star formation rate and massive star formation efficiency of GMCs in our catalog by taking advantage of a complete \\CS survey toward IRAS point-like sources with FIR colors characteristic of UC HII regions \\citep{bron96} complemented by a new unpublished \\CS survey (L. Bronfman et al., in preparation). Hereinafter we refer to the whole \\CS dataset as \\CS survey and to the IRAS point-like UC HII regions as \\IRAS sources. We select massive star forming regions from the IRAS point source catalog, because most of the embedded massive star luminosity is emitted around $100 \\mu m$ in the Mid- and Far-Infrared (see Figure 2 in \\citet{faundez2004}). The IRAS point sources are selected by their FIR colors, using the \\citet{wood89b} criterion for UC HII regions, which trace well the population of young embedded massive stars \\citep{faundez2004} emitting also in the millimeter continuum. This continuum survey shows that massive cold cores, the very early stages of MSF, are associated to the central UC HII region. The \\CS emission is a great tracer in that it is only excited in regions of densities above  $10^{4}$ - $10^{5}$ cm$^{-3}$. We detected the line emission for about 75\\% of the candidates. This is a  complete sample of the most massive and luminous regions of massive star formation in the Galaxy. There are more recent Galactic surveys in the Mid-Infrared and Sub-Millimeter regions of the spectrum (GLIMPSE, MIPSGAL, and ATLASGAL), but (a) there is no complete line surveys to determine their kinematic information and (b) only a small fraction of the luminosity from massive star forming regions comes in the NIR  and sub-mm wavelengths \\citep{faundez2004,tanti2012}.\\\\  In Section \\ref{catalog}  we present for the first  time a complete catalog of molecular clouds in the IV quadrant from $l = $ 300\\deg\\textrm{ }to $l = $ 348\\deg, within the solar circle. Distances, masses, and other physical properties of these objects are determined, and their statistical distributions are discussed. In Section \\ref{spiral}, the large scale spiral structure traced by GMCs within the Galactic disk is analyzed. In Section \\ref{msfr} we study the massive star formation rate per unit \\HH mass (MSFR) and massive star formation efficiency ($\\epsilon$) for GMCs. In Section \\ref{concl} the main conclusions of the present work are summarized. In Appendix \\ref{app1} the effect of subtracting a model from the CO data is described, and in Appendix \\ref{app2}, the two-fold distance ambiguity resolution for GMCs in our catalog is explained.\\\\ ",
        "conclusions": "}  Using the Columbia - U. de Chile CO Survey of the Southern Milky Way we present the first catalog of GMCs in the fourth Galactic quadrant, within the solar circle. Their average size, linewidth, and mass are estimated and their statistical properties are determined.\\\\  \\textit{GMCs Identification and Physical Properties}\\\\ We identify 92 molecular complexes in the inner Southern Galaxy and remove the two-fold distance ambiguity for 87 of them. The total molecular mass of GMCs is \\MHH $=$ 1.14 $\\pm$ 0.05 $\\times$ 10$^{8}$ M$\\solar$. The evidence in our catalog from the observational size-to-linewidth (\\dv $ \\propto R^{\\alpha}$) and size-to-\\HH volume density (\\nh $ \\propto R^{-\\beta}$)  relationships suggests that GMCs are close to virial equilibrium ($\\alpha + \\beta/2 =$ 0.95). The mass spectrum of the clouds (slope $\\gamma =$ 1.50 $\\pm$ 0.40) shows that most of the molecular mass in concentrated to the high mass end of the distribution. From the total molecular mass in our sample, 60\\% is located at the near distance while the remaining 40\\% is at the far distance. The completeness limit of the molecular mass spectrum is at 75\\% of the total mass.\\\\  \\textit{Total GMCs Mass and Axisymmetric Model}\\\\ The total molecular mass of GMCs in our catalog accounts for 40\\% of the molecular mass derived from the axisymmetric analysis of the \\HH volume density in the fourth Galactic quadrant, while in previous studies it was only 20\\%. From the virial analysis, we estimate an average \\COH conversion factor GMCs $\\chi_{GMCs} = 1.71 \\pm 0.09 \\times 10^{20}$ [K \\kms]$^{-1}$ cm$^{-2}$ which is consistent with the Galactic average value $\\chi = 1.56 \\pm 0.05 \\times 10^{20}$ [K \\kms]$^{-1}$ cm$^{-2}$ \\citep{hunter97}. The average Galactic value of the conversion factor in \\citet{hunter97} is estimated for both, cloud and inter-cloud medium simultaneously. A higher conversion factor for GMCs and a lower conversion factor for the inter-cloud medium could also reproduce the estimated average Galactic value and increase the molecular mass of GMCs reducing the difference even further.\\\\  \\textit{Large Scale Spiral Structure}\\\\ We trace here three large scale spiral arms within the solar circle: the \\emph{Centaurus}, \\emph{Norma}, and \\emph{3-kpc expanding} arms. The  \\emph{Carina} arm is not traced because of its wide latitude. After fitting a logarithmic spiral arm model for each of the arms, tangent directions are found to be at roughly 310\\deg, 330\\deg, and 338\\deg\\textrm{ }respectively. These values are consistent with previous determinations of the tangent directions done by \\citet{alvarez90} and \\citet{bron92}. The \\emph{Centaurus} arm is the most prominent feature in the southern Galaxy. Its pitch angle $p =$ 13\\deg.4\\textrm{ }is large compare to the pitch angles of the \\emph{Norma} ($p =$ 6\\deg.6) and \\emph{3-kpc} ($p =$ 5\\deg.6) arms. The spatial maps of the arms show that the molecular gas contained in GMCs and UC HII regions (\\IRAS sources) are tightly correlated, being the \\emph{Norma} spiral arm the feature with the most massive clouds in our sample. The positions of the far clouds in the \\emph{3-kpc expanding} arm converge to the location of the far side of the arm traced in \\citet{dame2008} down to $l = $ 348\\deg.\\\\  \\textit{Massive Star Formation Rate for GMCs}\\\\ Massive stars form in GMCs. Making use of the \\CS survey \\citep{bron96} toward IRAS point-like sources with FIR colors characteristic of UC HII regions \\citep{wood89b} and a new unpublished \\CS survey, we analyzed massive star formation GMCs. The total molecular mass of GMCs with massive star formation activity accounts for 85\\% of the total molecular mass in our sample, and 60\\% of them have masses larger than 6 $\\times$ 10$^{6}$ M$\\solar$. From the spiral features identified in our catalog, the \\emph{Norma} spiral arm is the most active in forming massive stars. The GMCs in this arm have a total FIR luminosity of 1.85 $\\times$ 10$^{7}$ L$\\solar$, equivalent to 50\\% of the total FIR luminosity in our sample. The GMC G337.750$+$0.000 is the most prominent massive star forming region in the inner southern Galaxy, with a FIR luminosity of 6.45 $\\times$ 10$^{6}$ L$\\solar$, equivalent to 35\\% of the total FIR luminosity in the arm. The observational evidence is consistent with the idea that, in general, the most massive clouds exhibit most of the massive star formation in the Galactic disk.\\\\  The massive star formation rate per unit \\HH mass (MSFR) and the massive star formation efficiency $\\epsilon$ are estimated for GMCs in our catalog. From the molecular mass of GMCs in our sample, the massive star formation rate is MSFR $=$ 0.41 $\\pm$ 0.06 L$\\solar$/M$\\solar$. A separate estimation for the most massive and most active massive star forming GMCs in our sample, GMC G328.250$+$0.375, GMC G331.500$-$0.125, GMC G336.875$+$0.125, and GMC G337.750$+$0.000 belonging to the \\emph{Norma} spiral arm, yields MSFR $=$ 0.58 $\\pm$ 0.09 L$\\solar$/M$\\solar$.\\\\  The MSFR in GMCs is about a factor two higher than the values derived for the Galactic plane $\\sim$ 0.21 L$\\solar$/M$\\solar$ between 0.4 $\\leq$ R/R$\\solar \\leq $ 0.8 and peak value $\\sim$ 0.31 L$\\solar$/M$\\solar$ for 0.5 $\\leq$ R/R$\\solar \\leq $ 0.6 (where part of the \\emph{Norma} arm is located) derived by \\citet{bron00} after a correction for Helium abundance is applied to their results. The discrepancy could be explained in terms of the molecular mass difference between GMCs in the present catalog and the molecular mass derived from the axisymmetric model analysis done by  \\citet{bron88}. We derive an average massive star formation efficiency $\\epsilon$ for GMCs of $\\sim$ 3\\% percent of their available molecular mass.\\\\"
    },
    "1402/1402.4676_arXiv.txt": {
        "abstract": "Galaxy groups form the environment of the majority of galaxies in the local Universe, and many host an extended hot intra-group medium whose radiative cooling appears to fuel, and be stabilised by, feedback from AGN in group-central galaxies. Unfortunately studies of the physical properties of groups and the influence of AGN on their member galaxies and gaseous haloes have been limited by a lack of reliable representative samples of groups in the local Universe. To address this problem, we have assembled the Complete Local-Volume Groups Sample (CLoGS), an optically-selected statistically-complete sample of 53 groups within 80 Mpc, which we aim to observe in both low-frequency radio  and X-ray wavebands. We here describe results from the first half of the sample, for which X-ray and radio observations are complete. Roughly 55\\% of the groups have group-scale X-ray halos, of which $\\sim$65\\% have cool cores a similar fraction to that found in galaxy clusters. While 25 of the 26 group central galaxies host radio AGN, among the X-ray bright groups only the cool core systems are found to support central jet sources. ",
        "introduction": "\\label{s:intro} Galaxy groups form the environment of the majority of galaxies in the local Universe. Many groups also host an extended hot intra-group medium (IGM) whose radiative cooling appears to fuel, and be stabilised by, AGN feedback from group-central galaxies. The group environment is conducive to galaxy evolution; low relative velocities and small galaxy separations drive mergers and tidal interactions, while motion through the IGM can produce gas compression or stripping. As such, groups are a critical environment in which to study the interplay between galaxy evolution, the development of the hot IGM, and AGN feedback. However, until recently, studies of groups have been hampered by the lack of reliable, statistically representative samples. Optically selected samples tend to become unreliable for low mass systems owing to the small number of detected group member galaxies. X-ray selection in the nearby Universe is dependent on the ROSAT All-Sky Survey, which is known to be biased toward centrally concentrated, cool core systems at low luminosities \\citep{Eckertetal11}. This bias can be illustrated by comparing the cool-core fractions of clusters measured from statistical samples \\citep[e.g., $\\sim$50\\%, ][]{Sandersonetal06} with the best available estimates from non-statistical samples of groups \\citep[$\\sim$85\\%,][]{Dongetal10}. Without a reliable, statistical sample of nearby groups we cannot know whether this difference indicates a change in the IGM heating balance between groups and clusters, or whether it is merely a product of bias.  To resolve this problem, we have created a Complete Local-Volume Groups Sample consisting of 53 groups within 80~Mpc. We initially draw groups from the Lyon Galaxy Group sample \\citep{Garcia93}, selecting systems with at least 4 members and at least one luminous early-type galaxy (L$_B>$3$\\times$10$^{10}$L$_\\odot$). Groups must also have Declination $>$-30$^\\circ$, to ensure visibility from GMRT and VLA. We then refine and expand group membership, and exclude the richest and poorest groups, which either correspond to known clusters, or lack the numbers of galaxies needed for accurate estimation of properties such as velocity dispersion and spiral fraction. Further details of the selection can be found at \\texttt{http://www.sr.bham.ac.uk/$\\sim$ejos/CLoGS.html}.  Our goal is then to observe these optically selected groups in the X-ray, since the presence of a hot X-ray emitting IGM would confirm them as gravitationally bound systems, and in the radio, to allow examination of the AGN and star formation in the member galaxies. At present, we have observed the full 53-group sample with GMRT, using dual-frequency 235/610~MHz observations with a typical on-source time of 3-4 hours. In the X-ray we have observed a high-richness subsample of 26 groups, itself statistically complete, with a limiting sensitivity of L$_{0.5-7~keV}>$1.2$\\times$10$^{42}$~erg~s$^{-1}$. ",
        "conclusions": ""
    },
    "1402/1402.2670_arXiv.txt": {
        "abstract": "\\noindent In light of the tension in cosmological constraints reported by the \\planck\\ team between their SZ-selected cluster counts and Cosmic Microwave Background (CMB) temperature anisotropies, we compare the \\planck\\ cluster mass estimates with robust, weak-lensing mass measurements from the {\\it Weighing the Giants} (WtG) project.  For the 22 clusters in common between the \\planck\\ cosmology sample and WtG, we find an overall mass ratio of $\\left< M_{Planck}/M_{\\rm WtG} \\right> = 0.688 \\pm 0.072$.  Extending the sample to clusters not used in the \\planck\\ cosmology analysis yields a consistent value of $\\left< M_{Planck}/M_{\\rm WtG} \\right> = 0.698 \\pm 0.062$ from 38 clusters in common.  Identifying the weak-lensing masses as proxies for the true cluster mass (on average), these ratios are $\\sim\\!\\!1.6\\sigma$ lower than the default bias factor of 0.8 assumed in the \\planck\\ cluster analysis.  Adopting the WtG weak-lensing-based mass calibration would substantially reduce the tension found between the \\planck\\ cluster count cosmology results and those from CMB temperature anisotropies, thereby dispensing of the need for ``new physics'' such as uncomfortably large neutrino masses (in the context of the measured \\planck\\ temperature anisotropies and other data).  We also find modest evidence (at 95 per cent confidence) for a mass dependence of the calibration ratio and discuss its potential origin in light of systematic uncertainties in the temperature calibration of the X-ray measurements used to calibrate the \\planck\\ cluster masses.  Our results exemplify the critical role that robust absolute mass calibration plays in cluster cosmology, and the invaluable role of accurate weak-lensing mass measurements in this regard. ",
        "introduction": "\\label{sect:intro}  The \\planck\\ satellite has recently provided new, precise cosmological constraints based on measurements of Cosmic Microwave Background (CMB) temperature anisotropies \\citep[][hereafter \\Pcos]{Planck13_16}, confirming that a spatially flat $\\Lambda$CDM model provides an excellent description of the observable Universe. However, an uncomfortable result of the \\planck\\ team's analysis is that for a spatially flat $\\Lambda$CDM model, their cosmological constraints derived from the number density of galaxy clusters detected with \\planck\\ through the Sunyaev-Zel'dovich (SZ) effect are in tension with the constraints from the CMB temperature power spectrum: in particular, the \\planck\\ cluster analysis \\citep[][hereafter \\Pccos]{Planck13_20} prefers a significantly lower value of the amplitude of the matter power spectrum at late times, $\\sigma_8$, compared to the value extrapolated from CMB temperature anisotropies.   The \\planck\\ team suggest two possible explanations for this tension. One is that the calibration of their cluster mass estimates, which are based on hydrostatic mass measurements derived from {\\it XMM-Newton} X-ray observations, is biased low. An important aspect of that analysis is the assumed value of the mass bias, $\\left< M_{Planck}/M_{\\rm true} \\right> = (1-b)$.  Expecting some level of non-thermal pressure support to be present in even relatively relaxed clusters, they present two sets of results, one which sets $(1-b)\\equiv0.8$, and one which marginalizes over the range $0.7 < (1-b) < 1.0$.  Reconciling the observed tension between \\planck\\ SZ cluster counts and the CMB power spectrum would, however, require an even lower ratio.  Alternatively, keeping the \\planck\\ mass calibration fixed at $(1-b) = 0.8$, the \\planck\\ team argue that their CMB and SZ-cluster data could be explained by a species-summed neutrino mass of $\\Sigma m_{\\nu} = (0.58 \\pm 0.20)$~eV ($2.8\\sigma$ significance departure from zero), a result that is in tension with the 95 per cent confidence upper limit of 0.23~eV derived from the \\planck\\ CMB analysis (\\Pcos) in combination with WMAP low-multipole polarization \\citep{blw13}, high-multipole temperature anisotropy from SPT and ACT \\citep{rsz12,dln14}, and the 6df, SDSS and BOSS BAO surveys \\citep{bbc11,pxe12,aab12}, and earlier studies using WMAP CMB data and independent X-ray and optical cluster measurements \\citep{mar10c,rvj10}.  Marginalizing over $0.7 < (1-b) < 1.0$, the evidence for massive neutrinos is weakened but still present, $\\Sigma m_{\\nu} = (0.40 \\pm 0.21)$~eV ($1.9\\sigma$).  In this letter, we investigate the reliability of the \\planck\\ cluster mass measurements by comparing them with robust, independently derived weak-lensing mass measurements from the \\WtG\\ (WtG) project \\citep{laa14, kla14,alk14}.  As has been discussed in the literature, systematic uncertainty in the absolute calibration of cluster masses is currently the most significant challenge facing cluster cosmology \\citep[WtG; see also][]{vkb09,mar10,rwr10,sta11,bhd13,aem11,lsst12}. Weak lensing provides our most promising method to calibrate the absolute masses of clusters since it measures the total mass directly, without relying on baryonic tracers, and is expected, from simulations, to be accurate (i.e. exhibit minimal bias in the mean). The WtG project targeted a subset of 51 clusters in catalogs based on the ROSAT All-Sky Survey \\citep[RASS,][]{tru93}. The overlap between those X-ray detected clusters and the \\planck\\ cluster catalog \\citep[see ][hereafter \\Pclust]{Planck13_29} is substantial: 38 WtG clusters are also in the \\planck\\ cluster sample, 22 of which were used in the \\planck\\ cluster count cosmology analysis. The WtG masses hence provide an excellent external, independent dataset to assess the calibration of the \\planck\\ cluster masses.  The masses and mass ratios quoted in this paper assume a flat $\\Lambda$CDM cosmological model with $\\Omega_{\\rm m} = 0.3$ and $H_0 = 100 \\,h\\, \\mbox{km/s/Mpc}$, where $h=0.7$. ",
        "conclusions": "In light of the reported tension between the cosmological constraints inferred from \\planck\\ SZ cluster counts and \\planck\\ CMB temperature anisotropies, we have compared the \\planck\\ team's cluster mass estimates with weak-lensing-based mass measurements determined by the \\WtG\\ project.  For the 22 clusters in common between the \\planck\\ cluster sample and WtG that are used in the \\planck\\ cluster cosmology analysis, we find the average mass ratio to be $\\left< M_{Planck}/M_{\\rm WtG} \\right> = (1-b) = 0.688 \\pm 0.072$.  For all 38 clusters in common between the two studies, we find $\\left< M_{Planck}/M_{\\rm WtG} \\right> = 0.698 \\pm 0.062$.  These values are $\\sim1.6\\sigma$ below the default bias adopted by \\Pccos, $(1-b)=0.8$. Also the range over which they marginalize for their more conservative analysis (0.7--1.0) is significantly skewed compared to the WtG mass calibration -- half of the WtG range is not included. Adopting the WtG mass calibration would raise the value of $\\sigma_8$ inferred from the \\planck\\ SZ cluster counts and alleviate the tension with the Planck CMB results.  We find modest (95 per cent) evidence for a mass-dependence of the mass calibration, with a best-fit power-law of $M_{Planck} \\propto {M_{\\rm WtG}}^{0.68^{+0.15}_{-0.11}}$.  As a possible origin for such a mass dependence we identify the temperature-dependent calibration uncertainty of the X-ray hydrostatic measurements used to calibrate the \\planck\\ cluster mass estimates.  Weak-lensing measurements provide an excellent complement to X-ray and SZ data in enabling the robust calibration of cluster masses. In forthcoming work (Mantz et al., in prep.), we will present cosmological constraints utilizing the full WtG sample in conjunction with X-ray selected cluster catalogs and extensive Chandra follow-up data."
    },
    "1402/1402.7325_arXiv.txt": {
        "abstract": "In the latest ten years, a big effort has given to study the morphology and kinematics of polar ring galaxies: many steps forward and new discoveries on the structure and formation mechanisms for such systems have been made during this time thanks to high resolution photometric and spectroscopic data. In this paper, I review the latest results obtained for this class of galaxies, from both observational and theoretical studies. I focus on the analysis of the observed properties (e.g., structure, colours, age, metallicity, and kinematics) for narrow and wide polar ring galaxies. In particular, I compare AM~2020-504 and NGC~4650A, which are the two prototypes for narrow and wide polar rings, respectively. I discuss similarities and differences between the two kinds of systems and how they reconcile with the main formation scenarios proposed for this class of galaxies. ",
        "introduction": "Polar ring galaxies (PRGs) are multi-spin systems composed by a central spheroidal component, the host galaxy (HG), and a polar structure which orbits in a nearly perpendicular plane to the equatorial one of the HG. In 1990, Whitmore and collaborators compiled the Catalog of Polar Ring Galaxies, Candidates and Related Objects (PRC), which includes 157 galaxies, where 106 have a well-defined PRG morphology, i.e., the two components (HG and polar structure) are clearly detectable and/or kinematically decoupled. In the rest of the sample, galaxies show a perturbed structure, where the central object has polar or high-inclined feature. Recently, taking advantage of new large data set from surveys, the number of PRG candidates PRG grows up, also at higher redshifts than galaxies in the PRC \\citep[][see also Iodice et al., this volume]{Finkelman,Resh1997a}. In particular, \\citet{Moiseev2011} made a new catalogue of PRGs based on the Sloan Digital Sky Survey (SDSS), providing new 275 candidates for polar rings and related objects (see also Smirnova et al., this volume).  For all the included objects in the PRC, the polar structure is identified as a {\\it ring\\/} \\citet{Whitmore1990}. Only subsequent studies on the prototype of PRGs, NGC~4650A (Fig. \\ref{fig1}), have revealed for the first time that the polar structure in this object has the morphology and kinematics of a {\\it disk\\/}, rather than a ring \\citep[see][]{Arnaboldi1997, Iodice02, Gallagher2002, Swaters2003}. NGC~4650A does not seem the only polar disk galaxy in the universe, but other PRGs show similar characteristics, such as A0136-0801 (PRC A-1) and UGC~9796 (PRC A-6) in PRC and SPRC-27, SPRC-59, and SPRC-69 in the new SDSS-based Polar Ring Catalog \\citep[SPRC;][]{Moiseev2011}. Thus, the `PRG' morphological type currently includes both polar rings and polar disks. As I show in this review, they have a different structure (i.e., light distribution, colors, age, and star motions) and probably a different formation history (but this is still an open issue).  {\\it Why it is interesting to study PRGs?\\/} Given the unique geometry, as suggested by \\citet{Whitmore1990}, PRGs represent a new way to study galaxy structure and formation. The existence of two orthogonal components of the angular momentum makes the PRGs the ideal laboratory to derive the 3-dimensional shape of the potential (see Combes, this volume).  The multi-spin morphology can not be explained by the collapse of a single proto-galactic cloud, but some kind of interaction (galaxy-galaxy or galaxy-environment) need to be invoked in the formation history of these systems.  The gravitational interactions are the `carrying pillar' of the cold dark matter (CDM) model for galaxy formation \\citep[e.g.,][]{Col00}, and in this framework, both merger and gas accretion play a major role in building the structure of spheroid and disk (see Conselice, this volume). Thus, PRGs are among the best objects in the universe to study the physics of such processes.  In PRGs, as well as for all types of galaxies, the observed morphology, kinematics, structure, and physical properties (i.e., stellar population, gas and dust content, and metallicity) provide the record of the formation mechanism. They result from the relative contribution of each kind of interaction (i.e., merger, accretion, and tidal stripping), from the physical parameters (i.e., mass and gas ratios) and orbital configuration (i.e., trajectory and relative velocity) of the two (or more) interacting systems.  All the observed physical quantities listed above need to be reproduced by models of galaxy formation.  In this paper, I review the main observational properties of galaxies with narrow and wide polar rings/disks (Sect.~\\ref{obs}), which are the most reliable formation scenarios proposed for PRGs (Sect.~\\ref{theory}) and how well they reconcile with the observations (Sect.~\\ref{conf}). ",
        "conclusions": "In this paper I reviewed the latest results obtained for PRGs from both observational and theoretical studies. This class of galaxies includes different kind of morphologies, narrow rings, wide polar rings/disks, and related objects (e.g., multiple or low inclined rings). I focused on the analysis of the observed physical properties (e.g., structure, colours, age, metallicity, and kinematics) for narrow and wide PRGs. In particular, I compared AM~2020-504 and NGC~4650A, which are the two prototypes for narrow and wide PRGs, respectively.  I discussed similarities and differences between the two kinds of systems and how they reconcile with the main formation scenarios proposed for PRGs. The main results of this analysis are listed below.  \\begin{itemize}  \\item For both narrow and wide PRGs, total luminosity $L_B$ and \\ion{H}{I} mass $M_{\\rm HI}$ are uniformly distributed and vary in the ranges $8.8 \\times 10^{10} \\le L_B \\le 10.8 \\times 10^{10}$~L$_\\odot$ and $8 \\times 10^{10} \\le M_{\\rm HI} \\le 11 \\times 10^{10}$~M$_\\odot$. Both narrow and wide PRGs have, on average, a large amount of \\ion{H}{I}, that can be 2 or 3 times the total luminosity.  \\item The HG has similar (spheroidal) morphology, colours, and age in both AM~2020-504 and NGC~4650A, but very different kinematics. In particular, the profile of the stellar velocity dispersion is different. As a consequence, the two galaxies have a different position with respect to the FJ relation for ETGs (see Fig.~\\ref{fig4}, right panel). This is also observed for other few PRGs which seem to have a bimodal distribution in the $\\log{\\sigma}-M_B$ plane, with some HGs having a larger $\\sigma$ and few others a lower $\\sigma$ with respect to spheroids of comparable total luminosity. This bimodality seems to be independent from the morphological type (narrow or wide PRGs). This result needs to be confirmed with a larger sample of PRGs, since its implications on the formation scenarios for spheroids are quite relevant.  \\item Contrary to the HGs, the polar structures in narrow and wide PRGs have different morphology, baryonic mass, and kinematics, but, on average, similar oxygen abundance (Fig.~\\ref{fig3}). NGC~4650A and AM~2020-504 slso show different age and colours, being the narrow ring older and redder than the wide polar disk (see Tab.~\\ref{tab1}). Furthermore, the position of both types of PRGs with respect to the TF relation for spiral galaxies is on average the same, i.e., most PRGs lie on a parallel relation with respect to that for spirals, showing larger \\ion{H}{I} linewidths than expected for the observed total luminosity (see Fig.~\\ref{fig4}, left panel).  \\item The key physical parameters (i.e., baryonic mass, HG kinematics, and metallicity in the polar structure), that let to discriminate among the three formation scenarios suggested for PRGs (see Sect.~\\ref{theory}) were derived in several works for some PRGs, including NGC~4650A and AM~2020-504. By comparing them with the predictions from models, one can trace the formation history for each object. In particular, in the case of NGC~4650A the cold accretion of gas along a filament is successfully tested.  \\end{itemize}  All the above results suggest that there is not a unique formation scenario for PRGs, but different mechanisms need to be invoked to explain the whole range of the observed properties for this class of galaxies.  As already showed by \\citet{Bou03}, the analysis performed in this review confirms that both major merger and tidal accretion scenarios are able to account for the observed morphologies and kinematics of narrow as well as wide polar rings/disks.  The simulations of a cold accretion of pristine gas along a filament account for the formation of a polar disk, like NGC~4650A. In the metallicity - luminosity plane (see Fig.~\\ref{fig3}), there are other PRGs, both narrow and wide polar rings/disks, with similar or even lower oxygen abundance: if the low metallicity is a strong constraint in favour of the cold accretion, may polar rings (not disks!), form through this mechanisms?"
    },
    "1402/1402.6053_arXiv.txt": {
        "abstract": "Based on high-resolution spectra obtained with PHOENIX at Gemini-South, CRIRES at VLT-UT1, and ESPaDOnS at CFHT, we present new measurements of the radial and projected rotational velocities of 219 low-mass stars. The target likely membership was initially established using the Bayesian analysis tool recently presented in \\citet{2013malo}, taking into account only the position, proper motion and photometry of the stars to assess their membership probability. In the present study, we include radial velocity as an additional input to our analysis, and in doing so we confirm the high membership probability for 130 candidates: 27 in $\\beta$ Pictoris, 22 in Tucana-Horologium, 25 in Columba, 7 in Carina, 18 in Argus and 18 in AB Doradus and 13 with an ambiguous membership. Our analysis also confirms the membership of 57 stars proposed in the literature. A subsample of 16 candidates was observed at three or more epochs, allowing us to discover 6 new spectroscopic binaries. The fraction of binaries in our sample is 25\\%, consistent with values in the literature. Twenty percent of the stars in our sample show projected rotational velocities (\\vsini) higher than 30 \\kms\\ and therefore are considered as fast rotators. A parallax and other youth indicators are still needed to fully confirm the 130 highly probable candidates identified here as new {\\it bona fide} members. Finally, based on the X-ray emission of {\\it bona fide} and highly probable group members, we show that for low-mass stars in the 12-120 Myr age range, the X-ray luminosity is an excellent indicator of youth and better than the more traditionally used R$_{\\rm X}$ parameter, the ratio of X-ray to bolometric luminosity. ",
        "introduction": "The search for young low-mass stars in the solar neighborhood has gained tremendous momentum recently as the youth and intrinsic faintness of these objects make them prime targets for planet searches through high-contrast imaging. This is both because young planetary-mass companions shine brighter and because low-mass stars are less luminous, which loosens the imaging contrast requirements. Finding more Young Moving Group (YMG) members also has intrinsic interest: to refine these YMG characteristics (initial mass function, age, velocities, spatial extent), to complete the solar neighborhood census, and to improve and develop new youth indicators.  As YMG members share common kinematics and span similar age ranges, one can search for overlooked members through a detailed kinematic and photometric analysis of nearby stars. The Hipparcos mission, combined with multiple radial velocity studies (e.g., \\citealt{2011anderson,2012debruijne}) led to the identification of the most massive members of these groups, but the relatively shallow completeness magnitude ($V=7-9$) prevented the detection of low-mass members (i.e., later than M5). Assuming that YMG stars follow a typical Initial Mass Function (IMF), the vast majority of low-mass members are expected to have gone unnoticed. This led to a number of efforts to extend measurements of proper motion, radial velocity and parallax to the nearby low-mass dwarfs (e.g., \\citealt{2012shkolnik}).  Proper motion studies beyond the reach of $Hipparcos$ also benefitted greatly from the advent other all-sky surveys (e.g., \\citealt{2009lepine, 2013zacharias}) reaching fainter magnitudes. In particular, infrared proper motions can now be obtained by correlating the WISE and 2MASS \\citep{2012cutri} surveys, which is ideally suited to measure proper motions well into the substellar regime \\citep[][]{2013rodriguez,2013agagne}. These all-sky photometric surveys have now been supplemented by radial velocity (RV) measurements coming from various works. Surveys like the Radial Velocity Experiment \\citep[RAVE; ][]{2006steinmetz, 2008zwitter} and SDSS \\citep{2000york} performed RV measurements on a large sample of stars, while others were focussed on old low-mass dwarfs \\citep{1998delfosse,2009jenkins,2012chubak,2013soubiran}. These RV studies are key for studying Galactic dynamics. Finally, and not least, parallax measurements of solar neighborhood stars is spear-headed by the CTIOPI and RECONS group \\citep[][]{2014riedel,2011riedel, 2010riedel} and other works \\citep{2009lepine,2012shkolnik,2013liu} focusing on the lowest mass component of the solar neighborhood.  To complete the census of nearby young kinematic group members, kinematics alone is not enough to confirm membership, independent youth indicators are also mandatory to confirm the age of these stars. The high magnetic activity of young stars gives rise to various indicators such as $H\\alpha$ emission, X-ray and UV \\citep{2005preibisch, 2011shkolnik, 2011rodriguez, 2013rodriguez}. Other spectroscopic indicators include the Li abundance \\citep{2008mentuch,2009dasilva} and others like sodium and calcium doublets (Na; 8183-8195\\AA, Ca: 8498-8542\\AA) that trace low-surface gravity \\citep{2004lyo, 2006slesnick, 2008mentuch, 2009shkolnik}. Stellar rotation and gyrochronology can also be used for constraining the age of an isolated star through a comparison with young stars in open clusters of known ages such as the Hyades, Pleiades and IC2391 \\citep{2011irwin, 2009lopez, 2010messina,2010lopez}. Recently, \\citet{2012areiners} presented the angular momentum evolution in low-mass stars, which can also be used as a youth indicator. To assess the membership of a given star to the known YMGs, a complete description of these groups is essential: their space velocity, spatial extent, luminosity along with appropriate youth indicators. Based on such YMG descriptions, any given star candidate member of a YMG can have its observational properties (proper motion, radial velocity, absolute magnitude) compared against actual observations.  In the present paper, we present the follow-up of the 130 highly candidate members of young kinematic groups identified or re-identified in a previous paper \\citep[][ hereafter Paper I]{2013malo}. In that paper, we presented an analysis tool (BANYAN\\footnote{Bayesian Analysis for Nearby Young AssociatioNs: www.astro.umontreal.ca/$\\sim$malo}) based on Bayesian inference to assess the membership probability of a given low-mass star to nearby YMGs. This analysis has unveiled a large population of potentially new young low-mass stars but those stars cannot be confirmed as {\\it bona fide} members of their respective association until knowledge of their radial velocity, parallax measurements and signs of youth are secured. The present work focuses on the radial and projected rotational velocity measurements which were measured through the cross-correlation technique using the infrared and optical domain. The paper is structured as follows. The sample and the observations are presented in $\\S$~\\ref{chap:deux} and data reduction is presented in $\\S$~\\ref{chap:trois}. A presentation of the radial and projected rotational velocity measurements follows in $\\S$~\\ref{chap:quatre}. The identification and confirmation of highly candidate members is described in $\\S$~\\ref{chap:cinq} and $\\S$~\\ref{chap:six} is devoted to a discussion of rotation-age and X-ray luminosity relations. Concluding remarks follow in $\\S$~\\ref{chap:sept}. ",
        "conclusions": "\\label{chap:sept}  The study aims at extending the census of low-mass star members of seven young nearby associations. A Bayesian statistical analysis described in Paper I was developed for identifying candidates based on a minimal set of observables (proper motion, photometry, position in the sky) combined with a kinematic model of the association. This method yields a membership probability to a given association as well as the most probable (statistical) distance and the predicted radial velocity. In the present study, radial velocity measurements were used to better constrain the membership of the candidates.  Starting from a sample of 920 stars, all showing indicators of youth such as H$\\alpha$ and/or X-rays emission, our analysis selected 247 candidates with a membership probability over 90\\%. To confirm the predicted membership, we have secured radial velocity measurements for 202 candidate members. These measurements were combined with a compilation of 166 measurements from the literature to select a list of 130 young K and M stars with a membership probability to all YMGs exceeding 90\\% when the radial velocity is included in our Bayesian analysis. A subgroup of 117 candidates are associated with a single YMG distributed as follows : 27 in $\\beta$PMG, 22 in THA, 25 in COL, 7 in CAR, 18 in ARG and 18 in ABDMG.  We investigated the rotation-age relation and found that stellar rotation (\\vsini) of young M dwarfs is significantly higher compared the their field counterparts. We also find that the X-ray luminosity of $\\sim$12-22 Myr-old $\\beta$PMG members are typically two orders of magnitude more luminous compared to field stars; this luminosity excess is a factor of $\\sim$4 compared to older ($\\sim$100 Myr) ABDMG members. Thus, the X-ray luminosity appears to be an excellent age discriminant for M dwarfs.  This work has unveiled a large population of highly probable low-mass members to nearby YMGs. A parallax, and ideally other indications of youth, are mandatory to firmly establish them as {\\it bona fide} members of their respective association."
    },
    "1402/1402.1498_arXiv.txt": {
        "abstract": "We study dwarf satellite galaxy quenching using observations from the Geha et al. (2012) NSA/SDSS catalog together with $\\Lambda$CDM cosmological simulations to facilitate selection and interpretation.  We show that fewer than $30 \\%$ of dwarfs ($M_\\star \\simeq 10^{8.5-9.5} \\msun$) identified as satellites within massive host halos ($M_{\\rm host} \\simeq 10^{12.5 - 14} \\msun$) are quenched, in spite of the expectation from simulations that half of them should have been accreted more than 6 Gyr ago. We conclude that whatever the action triggering environmental quenching of dwarf satellites, the process must be highly inefficient.  We investigate a series of simple, one-parameter quenching models in order to understand what is required to explain the low quenched fraction and conclude that either the quenching timescale is very long ($> 9.5 \\gyr$, a ``slow starvation'' scenario) or that the environmental trigger is not well matched to accretion within the virial volume.  We discuss these results in light of the fact that {\\em most} of the low mass dwarf satellites in the Local Group are quenched, a seeming contradiction that could point to a characteristic mass scale for satellite quenching. ",
        "introduction": "\\label{sec:intro}   The cessation of star formation in galaxies is a well studied and yet poorly understood phenomenon. That galaxies, when placed on a color-magnitude diagram, can roughly be divided into a red, non-star-forming sequence of mostly ellipticals and a blue, star-forming cloud of mostly spirals has been observed at least out to $z \\sim 1$ \\citep{Baldry:2004cd, Bell:2004vn, Borch2006, Cooper2007}. It is also observed that while the number density of star-forming galaxies is relatively constant at $z < 1$, non-star-forming (or quenched) galaxies have roughly doubled in number density over the past $7$--$10\\,\\gyr$ \\citep{Bell:2004vn, Borch2006, Faber:2007ef, Bundy:2007gh, Brown:2007ij}.  Environment appears to play a significant role in star formation quenching, and it may be the dominant driver at the lowest mass scales \\citep{Cooper:2006fk, Weinmann:2006fk, Capak:2007kx, Peng:2010dq, Geha:2012kx, Woo:2013w}.  Studies comparing satellite galaxies to isolated field systems of similar stellar mass in the local Universe find that satellites tend to exhibit lower star-formation rates, more bulge-dominated morphologies, as well as older and more metal-rich stellar populations \\citep{Baldry:2006b, vandenBosch:2008v, Yang:2007, Cooper2010, Pasquali:2010p}.  This observed suppression of star formation in satellite galaxies is commonly referred to as ``environmental quenching''.  Proposed quenching mechanisms that operate preferentially on satellite galaxies within dense environments (such as groups and clusters) range from ``strangulation\" -- when the larger potential well of the host dark matter halo accretes all of the gas that would have originally fueled star formation in the satellite galaxy \\citep{Larson:1980ve, Kawata:2008ab}, to ``harassment\" -- by which close encounters between densely packed cluster or group members strip gas from around the interacting galaxies \\citep{Moore:1996kl}, to ``ram-pressure stripping\" -- where the cold dense gas at the center of the satellite is violently removed from the galaxy as a result of a high-speed interaction with the hot gas halo of the host \\citep{Gunn:1972g, Bekki:2009b}.  Determining exactly how environmental quenching proceeds is complicated by the fact that some galaxies in the field are quenched. For isolated galaxies, free of the influence of high-density environments (i.e.~a more massive parent halo), the quenched fraction increases significantly with stellar mass \\citep{Peng:2010dq, Weinmann:2006fk}. This suggests that galaxies, at least massive galaxies, quench their own star formation independent of their environment. These observational results are supported by hydrodynamical simulations of isolated systems, in which star formation is suppressed owing to the formation of a merger-induced stellar spheroid that stabilizes the disk to the collapse of molecular clouds \\citep{Dekel:2009, Martig:2009ly}. This possible quenching mechanism is unique in that it depends on the mass of the stellar spheroid rather than that of the parent dark matter halo (i.e.~environment). Owing to the fact that self- and environmental quenching primarily affect the quenched fraction at different mass scales (with environmental quenching playing a primary role in low stellar mass systems, and mass quenching primarily affecting galaxies with a higher stellar mass), several authors have argued for a separation of quenching at low redshift into two functions -- one of environment and one of stellar mass \\citep{Kovac2014, Peng:2010dq}. Disentangling these two potential quenching regimes and identifying the associated physical mechanisms that operate within them remains a challenge for modern studies of galaxy evolution at low and intermediate redshift.  In overcoming these challenges, it is therefore desirable to model {\\em both} of these quenching mechanisms. Several studies have attempted to model the effects of both self- and environmental quenching using semi-analytic prescriptions. The resulting models typically associate quenching with various properties of a galaxy, such as the time since it fell into its host, in order to match observed quenched fractions as a function of stellar mass while also reproducing the observed stellar mass function \\citep{Wetzel2013, Cohn2013}. However, this overlapping of effects at stellar masses greater than $\\sim 10^{10}~\\msun$ continues to make a separate analysis of each mechanism difficult.  Dwarf galaxies offer a particularly useful avenue for disentangling the effects of environmental and self-quenching. As shown by \\citet[hereafter G12]{Geha:2012kx}, effectively all isolated dwarf galaxies are star-forming, suggesting that self-quenching for these systems is negligible. Specifically, G12 found that non-star-forming (or passive) galaxies with stellar mass below $\\sim 10^9\\,\\msun$ are absent beyond $\\sim 1$--$1.5\\,\\mpc$ of a luminous neighbor in the local Universe \\citep[see also][]{Wang2009}. The observations suggest that dwarf galaxies at these low masses do not shut down their own star formation; instead quenching requires the presence of a more massive neighbor. These galaxies therefore provide an ideal experimental sample: they do not self-quench, so any observed quenching can be directly related to environmental influences.  In addition to isolating the impact of environment on low-mass galaxies, the analysis by G12 provides well-defined, quantitative observational constraints on potential quenching mechanisms. For example, G12 probe the physical extent of environmental quenching, finding that dwarf galaxies are only quenched within a projected distance of $\\sim 1$--$1.5\\,\\mpc$ of a luminous galaxy. They interpret this result by suggesting that dwarfs are being quenched at $2$--$4$ times the virial radii of their Milky Way size hosts.  Below, we use cosmological simulations to suggest an alternative scenario in which the $\\sim 1-1.5 \\, \\mpc$ scale more likely corresponds to the virial radius of galaxy groups or small clusters.  Moreover, G12 also find that even within $250\\,\\kpc$ of a luminous neighbor, the fraction of quenched systems peaks at $\\sim 25$--$30\\%$. The implication of these findings is that satellite quenching is remarkably inefficient at low stellar masses, and we can constrain quenching models through the requirement that they match this low efficiency.  In what follows we compare mock observations of a large $N$-body simulation to the observed quenched fractions vs. projected distance to a luminous neighbor from G12. We consider several one-parameter models for quenching, including a simple quenched-at-infall scenario as well as models where quenching depends on the host $\\vmax$, satellite infall time, or the ratio of the satellite's current $\\vmax$ to the $\\vmax$ it had at infall. In Sections \\ref{sec:obs} and \\ref{sec:sim} we describe the G12 observational sample and our simulations, respectively.  Section \\ref{sec:fquench} presents our principal results.  Finally, we summarize our findings in Section \\ref{sec:summ} and discuss them in light of past work in Section \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:summ}  We have used mock observations of the \\msii~simulation \\citep{Boylan-Kolchin:2009ly} in order to interpret dwarf galaxy ($\\mstar \\simeq 10^9\\,\\msun$) quenching as a function of projected distance to a more luminous neighbor as observed by \\citet[G12]{Geha:2012kx}. Dwarfs of this mass are particularly useful as a test bed for environmental quenching because, as shown by G12, they are never (or almost never) quenched in isolation.  We show that while dwarfs within $250\\,\\kpc$ and $1000 \\,\\kms$ of a luminous neighbor are subhalos of a larger host approximately $90\\%$ of the time, only about $25$--$30\\%$ of such dwarfs are quenched (suggesting that $\\sim 30\\%$ of subhalos are quenched at these masses). The implication is that whatever is giving rise to subhalo quenching, the process must be fairly inefficient, at least when evaluated relative to the subhalo population as a whole.  We investigate a model in which dwarfs become quenched only after a time $\\tau_{\\rm infall}$ of being accreted into a larger host.  The required quenching timescale is quite long, $9.5\\,\\gyr$ (cf. Figure \\ref{fig:tinf}), compared to some estimates in the literature (see \\S \\ref{sec:disc}).  Alternatively, if dwarf quenching is instead related to the tidal forces experienced by the subhalo, then a simple model in which dwarfs with $\\vmax/\\vinfall < 0.65$ become quenched does a reasonable job in reproducing the observed quenched fraction in the innermost bin of G12 (see Figure \\ref{fig:vrat}), though in detail the infall time model provides a better match overall (Figure \\ref{fig:tinf}). If, instead, we try to explain the relatively low quenched fraction for subhalos by demanding that only hosts larger than a critical $\\vmax$ are able to quench their satellites, we find that a threshold $\\vmax$ of $650\\,\\kms$ is required. This value is unreasonably high, as it is well above that of isolated galaxy size hosts that are known to quench at least some of their satellites \\citep{More:2010kx,Monachesi2011,Tollerud:2011,Phillips2014}.  Quenching is likely a complicated process that depends on more than a single parameter. Our aim in this work is to focus on simple models in order to gain qualitative insight, though we also explore slightly more complicated cases in which two conditions must be met before quenching occurs. Combining a minimum host $\\vmax$ with a maximum ratio of satellite $\\vmax / \\vinfall$ does not significantly change the required values for either parameter. A minimum host $\\vmax$ of $350\\,\\kms$ only moves the best fit threshold ratio to $\\vmax / \\vinfall = 0.70$, which is not significantly larger than the required one-parameter model value of $0.65$. We also investigate a model that requires both a minimum host $\\vmax$ as well as a minimum $\\rm \\tau_{infall}$ for a satellite to become quenched. However, even requiring the relatively high host $\\vmax$ lower limit of $350~\\kms$ does not reduce the required infall time significantly. We estimate that these joint criterion for quenching will lead to a reduced threshold quenching timescale of $\\sim 7.5~\\gyr$.  Finally, combining $\\vmax / \\vinfall$ and $\\rm \\tau_{infall}$ does not significantly change the required values for either parameter. This is due to the large amount of overlap between subhalos that have lost a certain fraction of their $\\vinfall$ and those that have been orbiting within the virial radius of their hosts for a minimum amount of time."
    },
    "1402/1402.1983_arXiv.txt": {
        "abstract": "{{ } } { The primordial power spectrum describes the initial perturbations in the Universe which eventually grew into the large-scale structure we observe today, and thereby provides an indirect probe of inflation or other structure-formation mechanisms. Here, we introduce a new method to estimate this spectrum from the empirical power spectrum of cosmic microwave background (CMB) maps. } { A sparsity-based linear inversion method, coined \\textbf{PRISM}, is presented. This technique leverages a sparsity prior on features in the primordial power spectrum in a wavelet basis to regularise the inverse problem. This non-parametric approach does not assume a strong prior on the shape of the primordial power spectrum, yet is able to correctly reconstruct its global shape as well as localised features. These advantages make this method robust for detecting deviations from the currently favoured scale-invariant spectrum. } { We investigate the strength of this method on a set of WMAP 9-year simulated data for three types of primordial power spectra: a nearly scale-invariant spectrum, a spectrum with a small running of the spectral index, and a spectrum with a localised feature. This technique proves to easily detect deviations from a pure scale-invariant power spectrum and is suitable for distinguishing between simple models of the inflation. We process the WMAP 9-year data and find no significant departure from a nearly scale-invariant power spectrum with the spectral index $n_s = 0.972$. } { A high resolution primordial power spectrum can be reconstructed with this technique, where any strong local deviations or small global deviations from a pure scale-invariant spectrum can easily be detected. } ",
        "introduction": "\\label{sec:intro}  The primordial power spectrum encodes the physics of the early Universe and its measurement is one of the key research areas in modern cosmology. Amongst the proposed models which describe the early Universe, inflation \\citep{guth,Linde-inflation1982} is currently the most favoured one. In this model early perturbations are produced by quantum fluctuations during the epoch of an accelerated expansion. These perturbations then grow into the large scale structure we observe today. The simplest models of inflation predict almost purely adiabatic primordial perturbations with a nearly scale-invariant power spectrum. In these models the primordial power spectrum is often described in terms of a spectral index $n_s$ and an amplitude of the perturbations $A_s$ as \\begin{equation} P(k)=A_s \\left(\\frac{k}{k_p}\\right)^{n_s-1}\\;, \\label{eq:pkdef} \\end{equation} where $k_p$ is a pivot scale. This spectrum represents the initial conditions set at inflation. The simplest ansatz for characterising the primordial perturbations is the so-called Harrison-Zeldovich (HZ) model, which sets $n_s=1$ \\citep{Harrison1970,Zeldovich1972}. This is an exact scale-invariant spectrum, which has been ruled out by different datasets, as will be discussed later. Instead, the near scale-invariant spectrum with $n_s<1$ fits the current observations very well \\citep[for e.g.,][]{PlanckCP}. However, numerous models have been proposed for the generation of the perturbations, predicting deviations from the perfectly scale-invariant power spectrum. The simplest are the slow-roll inflationary models which describe the deviations through a minimal scale dependence of the spectral index of the power spectrum, the so-called `running' $\\alpha_s$, formulated as \\begin{equation} P(k)=A_s \\left(\\frac{k}{k_p}\\right)^{n_s-1+\\frac{1}{2}\\alpha_s\\ln \\left({k}/{k_p}\\right)}\\;. \\label{eq:pkdef} \\end{equation} More complex models generating deviations from scale-invariance include those with features on the potential \\citep{Starobinsky:1992ts,Adams:2001vc,Wang:2002hf,Hunt:2004vt,Joy:2007na,Hunt:2007dn,Pahud:2008ae,Lerner:2008ad,Kumazaki2011,Meerburg2012}, a small number of $e$-folds \\citep{Contaldi:2003zv,Powell:2006yg,Nicholson:2007by}, or other exotic inflationary models \\citep{Lesgourgues:1999uc,Feng:2003zua,Mathews:2004vu,Jain:2008dw,Romano:2008rr}. Therefore, determining the shape of the primordial power spectrum will allow us to evaluate how well these models of the early Universe compare to the observations, rule out some of the proposed models, and thus giving us a better intuition into the conditions of the primordial Universe.  A few probes of the physics of the early Universe include non-Gaussianity, the primordial tensor power spectrum, a cosmic gravitational wave background and a cosmic neutrino background, none of which have been observed with an acceptable significance. On the other hand, we can observe $P(k)$ through the windows of the Cosmic Microwave Background (CMB) and Large Scale Structure (LSS), which are incredibly important and powerful insights into the early Universe.  The recent Planck mission temperature anisotropy data, combined with the WMAP large-angle polarisation, constrain the scalar spectral index to $n_{s}=0.9603 \\pm0.0073$ \\citep{PlanckCP}, which rules out exact scale-invariance at over $5 \\sigma$. In addition, Planck does not find a statistically significant running of the scalar spectral index, obtaining $\\alpha_s =-0.0134\\pm0.0090$. In \\citet{PlanckCP} an extensive investigation is performed by the Planck collaboration to see whether the primordial power spectrum contains any features. They report that a penalised likelihood approach suggests a feature near the highest wavenumbers probed by Planck at an estimated significance of $\\sim 3 \\sigma$. In addition, a parameterised oscillatory feature does improve the fit to the data by $\\Delta\\chi^2_{\\mathrm{eff}} \\approx 10$, however Bayesian evidence does not prefer these models. On the other hand, high resolution CMB experiments, such as the South Pole telescope (SPT)\\footnote{\\url{http://pole.uchicago.edu/spt/index.php}}, detect a small running of the spectral index; $-0.046 < \\alpha_s < -0.003$ at $95\\%$ confidence \\citep{SPT_cos}. In general, any detections of the running of the spectral index have been small and consistent with zero. Therefore, a highly sensitive algorithm is required to detect these small deviations.  There are generally two approaches to determine the shape of the primordial power spectrum, one is by parametrisation and the second is a reconstruction. Numerous parametric approaches search for features with a similar form to those in complex inflationary models have been performed along with a simple binning of $P(k)$ \\citep{Bridle:2003sa,Contaldi:2003zv,Parkinson:2004yx,Sinha:2005mn,Sealfon:2005em,Mukherjee:2005dc,Bridges:2005br,Bridges:2006zm,Covi:2006ci,Joy:2008qd,Verde:2008zz,PaykJaffe_OB,GuoEtAl2011,GoswamiPrasad2013}. Non-parametric methods, which make no assumptions about the model of the early Universe, have also been probed \\citep{Hannestad:2000pm,Wang:2000js,Matsumiya:2001xj,Shafieloo:2003gf,Bridle:2003sa,Kogo:2003yb,Mukherjee:2003cz,Mukherjee:2003ag,Hannestad:2003zs,Kogo:2004vt,TocchiniValentini:2004ht,Leach:2005av,Shafieloo:2006hs,Shafieloo:2007tk,Nagata:2008tk,Nagata:2008zj,Nicholson:2009pi,NichPaykCont,HazraEtAl20131}. For an extensive review on how to search for features in the primordial power spectrum using a wide range of methods, refer to the following papers and the references therein, which provide a sample on non-parametric reconstruction:  deconvolution \\citep{Tocchini-Valentini2006,IchikiNagata2009,IchikiEtAl2010} (including Richardson-Lucy deconvolution \\citep{Lucy, Richardson,Hamann2010, Shafieloo:2007tk}), smoothing splines \\citep{VP-mini-param-pPS,PeirisVerde2010, Sealfon:2005em,GauthierBucher2012}, linear interpolation \\citep{Hannestad:2003zs,Bridle:2003sa}, and Bayesian model selection \\citep{BridgesEtAl2009,VazquezEtAl2012}.  Non-parametric methods are hampered by the non-invertibility of the transfer function that descries the {\\it transfer} from $P(k)$ to CMB (or LSS). Specifically for the CMB power spectrum, the dependence on the transfer function has the form \\begin{equation} C_{\\ell}^{\\textrm{th}}=4\\pi\\int_{0}^{\\infty}d\\ln k\\Delta_{\\ell}^{2}(k)P(k)\\;, \\label{eq:CMB} \\end{equation} where $\\ell$ is the angular wavenumber that corresponds to an angular scale via $\\ell\\sim180^{o}/\\theta$ and $\\Delta_{\\ell}(k)$ is the angular transfer function of the radiation anisotropies, which holds the cosmological parameters responsible for the evolution of the Universe. As the CMB spectrum is \\textit{jointly} sensitive to the primordial spectrum and the cosmological parameters in the transfer function, there is an induced degeneracy between them. The impact and level of this degeneracy have been investigated in \\citep{PaykJaffe_OB}. A joint estimation of the cosmological parameters and a free form primordial power spectrum would be prohibitively expensive to perform (as the parameter space potentially becomes very large). As a result, a parametric form of the primordial power spectrum is assumed when jointly estimating this spectrum along with the other cosmological parameters. This hides any degeneracies between the cosmological parameters in the transfer function and the form of $P(k)$. Thus it is not clear what the significance of any features found in the reconstructed $P(k)$ should be. One way to break this induced degeneracy is by adding extra information, such as polarisation or LSS data \\citep{Hu:2003vp,Nicholson:2009pi,Mortonson:2009qv}.  The other hurdle into the estimation of the primordial spectrum is that this continuous spectrum is deconvolved from discrete data $C_{\\ell}$. This causes problems if the primordial power spectrum contains features that are smaller or comparable to the gridding in $\\ell$ ($\\Delta \\ell=1$). This limits our ability to fully recover the primordial power spectrum; in the case of the CMB, even a perfect survey cannot recover the primordial power spectrum completely \\citep{ho}.  Here, we propose a new non-parametric method for the reconstruction of the primordial power spectrum from CMB data, which is based on the sparsity of the primordial power spectra in a wavelet basis and an appropriate noise modelling of the CMB power spectrum \\citep{Tousi}.    \\subsection*{Paper content} \\label{sec:content}  In section~\\ref{sec:method} we present the primordial power spectrum reconstruction problem and describe the technique we have developed to perform the reconstruction. Our algorithm is tested on three sets of simulated spectra and applied to WMAP 9-year data in section~\\ref{sec:results}. In section~\\ref{sec:conclusion} we conclude and state some potential perspectives. ",
        "conclusions": "\\label{sec:conclusion}  The primordial power spectrum describes the initial perturbations in the Universe and hence provides an indirect probe of inflation or other structure-formation mechanisms. The simplest models of inflation are the most favoured by the data and predict a nearly scale-invariant power spectrum with a small running. One way to measure this spectrum is through the windows of the CMB data. The problem, though, is that the singular nature of the radiation transfer function and the joint estimation of the cosmological parameters in the transfer function and the primordial power spectrum, along with the different types of noise sources impose a limit into the full recovery of the primordial spectrum. Therefore, devising a technique which is sensitive enough to detect deviations from scale-invariance is the key to recover an accurate primordial power spectrum.  In this paper we have introduced a new non-parametric technique, coined PRISM, to recover the primordial power spectrum from masked noisy CMB data. This is a sparse recovery method, which uses the sparsity of the primordial power spectrum as well as an adapted modelling for the noise of the CMB power spectrum. This algorithm assumes no prior shape for the primordial spectrum and does not require a coarse binning of the power spectrum, making it sensitive to both global smooth features (e.g., running of the spectral index) as well as local sharp features (e.g., a bump or an oscillatory feature). Another advantage of this method is that, thanks to the clever modelling of the sample variance on the input angular power spectrum, the regularisation parameter can be specified in terms of a signal-to-noise significance level for the detection of features. These advantages make this technique very suitable for investigating different types of departures from scale-invariance in the primordial power spectrum, whether it is the running of the spectral index or some localised sharp features as predicted by some of the inflationary models.  We have investigated the strength of our proposed algorithm on a set of WMAP 9-year simulated data for three types of primordial power spectrum; a near scale-invariant spectrum, a spectrum with a small running of the spectral index, and a spectrum with a localised feature. We have shown that our algorithm can easily recover the three spectra with an excellent accuracy in the range $k\\sim0.001-0.1 \\mathrm{\\;Mpc}^{-1}$. In addition, the errors in the recovered spectra are small enough that the three types of primordial spectra can easily be distinguished in the range $k\\sim0.015-0.1 \\mathrm{\\;Mpc}^{-1}$. This technique has proved to easily detect small global and localised deviations from a pure scale-invariant power spectrum and is suitable for distinguishing between simple models of the inflation.  Using PRISM, we have reconstructed a primordial power spectrum from the LGMCA WMAP 9-year data and have investigated possible departures from the WMAP 9-year near scale-invariant spectrum. We have not detected any significant deviations from this simple model of the primordial power spectrum. We have demonstrated the feasibility of using PRISM on masked CMB data contaminated by instrumental noise. Better constraints will be obtained in future works by processing Planck data which provides a much lower instrumental noise, thus improving the range of scales we are able to probe with much better accuracy.  To this end, we also acknowledge previous algorithms aimed at reconstructing the primordial power spectrum with no need for binning, most of which have been referenced in this paper. The most recent work is by \\citet{HazraEtAl20131}, who use an adapted and improved Richardson-Lucy algorithm, dubbed MRL, to reconstruct the primordial power spectrum. Due to the very high level of the instrumental noise on small scales in the WMAP 9-year data, the MRL algorithm can only take the unbinned CMB spectrum for $\\ell<900$. For larger angular scales, $\\ell=900-1200$, a binned CMB spectrum is used. In addition, due to the induced artefacts in the reconstructed primordial spectrum, a smoothing step is necessary after the reconstruction is performed. Henceforth, compared to the MRL algorithm, the advantage of our algorithm is twofold. One is the ability to use the unbinned CMB spectrum for the whole multipole range $\\ell=2-1200$. This is because of our accurate noise modelling on the CMB power spectrum. In addition, unlike the MRL algorithm, there is no need to smooth the spectrum after the reconstruction as we look for the sparsest solution in our algorithm. This is why PRISM performs significantly better than previous algorithms, including the MRL algorithm.  The developed C++ and IDL codes will be released with the next version of iSAP (Interactive Sparse astronomical data Analysis Packages) via the web site\\\\ \\\\ {\\centerline{\\texttt{http://cosmostat.org/isap.html} .}}\\\\ \\\\ All results have been obtained using the isap routine \\textbf{mrs\\_prism} with the following command line: \\begin{verbatim} pk = mrs_prism(Cl, noise=Nl, TransferMat=Mat) \\end{verbatim} where \\texttt{Cl} contains the observed pseudo-power spectrum of the masked noisy CMB maps, \\texttt{Nl} is an estimate of the instrumental noise power spectrum $N_\\ell$, and \\texttt{Mat} is the input transfer matrix which include the effects of the radiation transfer function, the mask, the beam and the HEALPix window (Equation \\ref{eq:FinalT}). The transfer matrix can be been computed using the isap routine \\textbf{mrs\\_transfer\\_matrix} and by default the transfer matrix is derived from the WMAP 9-year best fit cosmology model."
    },
    "1402/1402.6337_arXiv.txt": {
        "abstract": "SN~2011fe was the nearest and best-observed type Ia supernova in a generation, and brought previous incomplete datasets into sharp contrast with the detailed new data. In retrospect, documenting spectroscopic behaviors of type Ia supernovae has been more often limited by sparse and incomplete temporal sampling than by consequences of signal-to-noise ratios, telluric features, or small sample sizes. As a result, type Ia supernovae have been primarily studied insofar as parameters discretized by relative epochs and incomplete temporal snapshots near maximum light. Here we discuss a necessary next step toward consistently modeling and directly measuring spectroscopic observables of type Ia supernova spectra. In addition, we analyze current spectroscopic data in the parameter space defined by empirical metrics, which will be relevant even after progenitors are observed and detailed models are refined. ",
        "introduction": "\\label{s:intro}  The transient nature of extragalactic type Ia supernovae (SN~Ia) prevent studies from conclusively singling out unobserved progenitor configurations \\citep{Roelofs08,WLi11a,Kilic13}. It remains fairly certain that the progenitor system of SN~Ia comprises at least one compact C$+$O white dwarf \\citep{Chandrasekhar57,Nugent11,Bloom12}. However, \\emph{how} the state of this primary star reaches a critical point of disruption continues to elude astronomers. This is particularly so given that less than $\\sim$ 15\\% of locally observed white dwarfs have a mass a few 0.1M$_{\\odot}$ greater than a solar mass; very few systems near the formal Chandrasekhar-mass limit\\footnote[5]{By ``formal'' we are referring to the mass limit that omits stellar rotation (see Appendix of \\citealt{Jeffery06}).}, M$_{Ch}$ $\\approx$ 1.38 M$_{\\odot}$ \\citep{Vennes99,Liebert05,Napiwotzki05, Partha07,Napiwotzki07}.  Thus far observational constraints of SN~Ia have been inconclusive in distinguishing between the following three separate theoretical considerations about possible progenitor scenarios. Along side perturbations in the critical mass limit or masses of the progenitors, e.g., from rotational support \\citep{Muller85,Yoon05,Chen09,Hachisu12a,Tornambe13} or variances of white dwarf (WD) populations \\citep{Kerkwijk10,Dan13}, the primary WD may reach the critical point by accretion of material from a low-mass, radially-confined secondary star \\citep{Whelan73,Nomoto77,Hayden10SDSS,Bianco11,Bloom12,Hachisu12a,Wheeler12,Mazzali13,Chen13}, and/or through one of several white dwarf merger scenarios with a close binary companion \\citep{Webbink84,Iben84,Paczynski85,Thompson11,BoWang13,Pakmor13}. In addition, the presence (or absence) of circumstellar material may not solely rule out particular progenitor systems as now both single- and double-degenerate systems are consistent with having polluted environments prior to the explosion \\citep{Shen13,Phillips13}.   Meanwhile, and within the context of a well-observed spectroscopically normal SN~2011fe, recent detailed models and spectrum synthesis along with SN Ia rates studies, a strong case for merging binaries as the progenitors of normal SN Ia has surfaced (c.f., \\citealt{Kerkwijk10,WLi11b,Blondin12,Chomiuk13,Dan13,Moll13,Maoz13,Johansson14}). However, because no progenitor system has \\emph{ever} been connected to any SN Ia, most observational constraints and trends are difficult to robustly impose on a standard model picture for even a single progenitor channel; the SN~Ia problem is yet to be confined for each SN Ia subtype.  As for restricting SN Ia subtypes to candidate progenitor systems: (i) observed ``jumps'' between mean properties of SN Ia subtypes signify potential differences of progenitors and/or explosion mechanisms, (ii) the dispersions of individual subtypes are thought to arise from various abundance, density, metallicity, and/or temperature enhancements of the original progenitor system's post-explosion ejecta tomography, and (iii) ``transitional-type'' SN Ia complicate the already similar overlap of observed SN Ia properties \\citep{Nugent95a,Lentz00,Benetti05,Branch09,Hoflich10,WangX12,WangX13,Dessart13models}. Moreover, our physical understanding of all observed SN~Ia subclasses remains based entirely on interpretations of idealized explosion models that are so far constrained and evaluated by ``goodness of fit'' comparisons to incomplete observations, particularly for SN~Ia spectra at all epochs.  By default, spectra have been a limiting factor of supernova studies due to associated observational consequences, e.g., impromptu transient targets, variable intrinsic peak luminosities, a sparsity of complete datasets in wavelength and time, insufficient signal-to-noise ratios, and the ever-present obstacle of spectroscopic line blending \\citep{Payne40}. Subsequently, two frequently relied upon empirical quantifiers of SN~Ia spectroscopic diversity have been the rate at which rest-frame 6100 \\AA\\ absorption minima shift redward vis-\\`{a}-vis projected Doppler velocities of the absorbing Si-rich material \\citep{Benetti05,WangX09Subtype} and absorption strength measurements (a.k.a. pseudo equivalent widths; pEWs) of several lines of interest (see \\citealt{Branch06,Hachinger06,Silverman12maxlight,Blondin12}). Together these classification schemes more-or-less describe the same events by two interconnected parameter spaces (i.e. flux and expansion velocities, \\citealt{Branch09,Foley11,Blondin12}) that are dependent on a multi-dimensional array of physical properties. Naturally, the necessary next step for supernova studies alike is the development of prescriptions for the physical diagnosis of spectroscopic behaviors (see \\S2.2 and \\citealt{Kerzendorf14}).  For those supernova events that \\emph{have} revealed the observed patterns of SN~Ia properties, the majority are termed ``Branch-normal'' \\citep{Branch93,WLi11b}, while others further away from the norm are historically said to be ``peculiar'' (e.g., SN~1991T, 1991bg; see \\citealt{Flipper97} and references therein). Although, many non-standard events have since obscured the boundaries between both normal and peculiar varieties of SN~Ia, such as SN~1999aa \\citep{Garavini04}, 2000cx \\citep{Chornock00,WLi01,Rudy02}, 2001ay \\citep{Krisciunas11}, 2002cx \\citep{WLi03}, 2003fg \\citep{Howell06,Jeffery06}, 2003hv \\citep{Leloudas09,Mazzali11}, 2004dt \\citep{Wang06,Altavilla07}, 2004eo \\citep{Pastorello07a}, 2005gj \\citep{Prieto07}, 2006bt \\citep{Foley10b}, 2007ax \\citep{Kasliwal08}, 2008ha \\citep{Foley09,Foley10a}, 2009ig \\citep{Foley12c,Marion13}, PTF10ops \\citep{Maguire11}, PTF11kx \\citep{Dilday12,Silverman13}, and 2012fr \\citep{Maund13,Childress13}.  The fact that certain subsets of normal SN~Ia constitute a near homogenous group of intrinsically bright events has led to their use as standardizable distance indicators \\citep{Kowal68,Elias81,Branch92,Riess99,Perlmutter99,Schmidt04,Mandel11,Maeda11,Sullivan11SNLS,Hicken12}. However, this same attribute of homogeneity remains the greatest challenge in the individual study of SN~Ia given that the time-evolving spectrum of a supernova is unique unto itself from the earliest to the latest epochs.  Because SN~Ia are invaluable tools for both cosmology and understanding progenitor populations, a multitude of large scale surveys, searches, and observing campaigns\\footnote[6]{e.g., The Automated Survey for SuperNovae (Assassin), The Backyard Observatory Supernova Search (BOSS), The Brazilian Supernova Search (BRASS), The Carnegie Supernova Project (CSP), The Catalina Real-Time Transient Survey (CRTS), The CHilean Automatic Supernovas sEarch (CHASE), The Dark Energy Survey (DES), The Equation of State: SupErNovae trace Cosmic Expansion (ESSENCE) Supernova Survey, The La Silla-QUEST Variability Survey (LSQ), Las Cumbres Observatory Global Telescope Network (LCOGT), The Lick Observatory Supernova Search (LOSS), The Mobile Astronomical System of the Telescope-Robots Supernova Search (MASTER), The Nearby Supernova Factory (SNfactory), The Optical Gravitational Lensing Experiment (OGLE-IV), The Palomar Transient Factory (PTF), The Panoramic Survey Telescope and Rapid Response System (Pan-STARRS), The Plaskett Spectroscopic Supernova Survey (PSSS), Public ESO Spectroscopic Survey of Transient Objects (PESSTO), The Puckett Observatory World Supernova Search, The ROTSE Supernova Verification Project (RSVP), The SDSS Supernova Survey, The Canada-France-Hawaii Telescope Legacy Survey Supernova Program (SNLS), The Southern inTermediate Redshift ESO Supernova Search (STRESS), The Texas Supernova Search (TSS); for more, see \\url{http://www.rochesterastronomy.org/snimages/snlinks.html}.} are continually being carried out with regularly improved precision. Subsequently, this build-up of competing resources has also resulted in an ever growing number of new and important discoveries, with less than complete information for each. In fact, with so many papers published each year on various aspects of SN~Ia, it can be difficult to keep track of new results and important developments, including the validity of past and present theoretical explosion simulations and their related observational interpretations (see \\citealt{Maoz13} for the latest).  Here we compile some of the discussions on spectroscopic properties of SN~Ia from the past decade of published works. In \\S2 we overview the most common means for studying SN~Ia: light curves (\\S2.1), spectra (\\S2.2), and detailed explosion models (\\S2.3). In particular, we overview how far the well-observed SN~2011fe has progressed the degree of confidence associated with reading highly blended SN~Ia spectra. Issues of SN~Ia diversity are discussed in \\S3. Next, in \\S4 we recall several SN~Ia that have made up the bulk of \\emph{recent} advances in uncovering the extent of their properties and peculiarities (see also the Appendix for a guide of some recent events). Finally, in \\S5 we summarize and conclude with some observational lessons of SN~2011fe. ",
        "conclusions": "\\label{s:summary}  Observations of a significant number of SN Ia during the last two decades have enabled us to document a larger expanse of their physical properties which is manifested through spectrophotometric diversity. While in general SN Ia have long been considered a homogeneous class, they do exhibit up to 3.5 mag variations in the peak luminosity, whereas ``normal'' SN Ia dispersions are $\\sim$1 mag, and constitute several marginally distinct subtypes \\citep{Blondin12,Scalzo12,Silverman13IaCSM,Foley13,Dessart13models}. Consequently, the use of normal SN Ia for cosmological purposes depends on empirical calibration methods (e.g., \\citealt{Bailey09}), where one of the most physically relevant methods is the use of the width-luminosity relation \\citep{Phillips93,Phillips99}.  Understanding the physics and origin of the width-luminosity-relationship of SN Ia light curves is an important aspect in the modeling of SN Ia \\citep{Khokhlov93,Lentz00,Timmes03,Nomoto03,Kasen07,Kasen09,Meng11,Blondin11a}. Brighter SN Ia often have broad light curves that decline slowly after peak brightness. Slightly less bright or dimmer SN Ia have narrower and relatively rapidly declining light curves. In addition, several SN Ia do not follow the width-luminosity-relationship (e.g., SN~2001ay, 2004dt, 2010jn, SCC, CL and SN~2002cx-like SN Ia), which reinforces the notion that a significant number of physically relevant factors influence the diversity of SN Ia overall (see \\citealt{WangX12,Baron12}).  Despite the ever increasing number of caught-early supernovae, our perspective on their general properties and individual peculiarities undergoes a continual convergence toward a set of predictive standards with which models must be seen to comply. The most recent observational example is that of SN~2012fr \\citep{Maund13,Childress13}, a normal/low-velocity-gradient SN Ia that has been added to the growing list of similar SN Ia that exhibit stark evidence for a distinctly separate region of ``high-velocity'' material ($>$16,000 km s$^{-1}$). While the origin of high velocity features in the spectra of SN Ia is not well understood, it is concurrent with polarization signatures in most cases which implies some amount of ejecta density asymmetries (e.g., \\citealt{Kasen03,WW08,Smith11, Maund13}). Furthermore, since understanding the temporal behavior of high velocity \\ion{Si}{2}/\\ion{Ca}{2} depends on knowing the same for the photospheric component, studies that focus on velocity gradients and potential velocity-plateaus of the photospheric component could make clearer the significance of the physical separation between these two regions of material (see \\citealt{Patat96,Kasen03,Tanaka08,Foley12c,Parrent12,Scalzo12,Childress13,Marion13}). However, it is at least certain that all viable models that encompass ``normal'' SN Ia conditions must account for the range of properties related to velocity evolution (see \\citealt{Blondin12} and references therein), the occasionally observed however potentially under-detected signatures of C$+$O material at both low and high velocities \\citep{Thomas07,Parrent11,Thomas11,Folatelli12,Silverman12b,Piro13a,Mazzali13}, a high-velocity region of either clumps or an amorphous plumage of opaque Si-, Ca-based material \\citep{Gamezo04,Leonard05,Wang07,Maund10a,Piro11}, and the supposed blue/red-shift of nebular lines emitted from the inner IPE-rich material \\citep{Maedanature,McClelland13,Silverman13latetime}.  For at least normal SN Ia, there remain two viable explosion channels (with a few sub- and super-M$_{Ch}$ sub-channels) regardless of the hierarchical dominance of each at various redshifts and/or ages of galactic constituents (c.f., \\citealt{Roepke12,Hachisu12a,Seitenzahl13,Pakmor13,Moll13,Claeys14}). Also, it may or may not be the case that some SN~Ia are 2$+$ subtypes viewed upon from various lines of sight amidst variable CSM interaction \\citep{Maedanature,Foley12dust,Scalzo12,Leloudas13,Dan13,Moll13,Dessart13models}. However, with the current lack of complete observational coverage in wavelength, time, and mode (i.e. spectrophotometric and spectropolarimetric observations) for all SN~Ia subtypes and ``well-observed'' events, there is a limit for how much constraint can be placed on many of the proposed explosion models and progenitor scenarios. That is, despite observational indications for and theoretical consistencies with the supposition of multiple progenitor channels, the observed diversity of SN~Ia does not yet necessitate that each spectrophotometric subtype be from a distinctly separate explosive binary scenario than that of others within the SN~Ia family of observed events; particularly so for normal SN~Ia.  For the purposes of testing the multifaceted predictions of theoretical explosion models, time series spectroscopic observations of SN~Ia serve to visualize the post-explosion material of an unknown progenitor system. For example during the summer of 2011 astronomers bore witness to SN~2011fe, the best observed normal type Ia supernova of the modern era. The prompt discovery and follow-up of this nearby event uniquely allowed for a more complete record of observed properties than all previous well-observed events. More specifically, the full range (in wavelength and time) of rapid spectroscopic changes was documented with continual day-to-day follow-up into the object's post-maximum light phases and well beyond. However, the observational side of visualizing other SN Ia remains inefficient without the logistical coordination of many telescope networks (e.g., LCOGT; \\citealt{LCOGT}), telescopes large enough to make nearly all SN ``nearby'' in terms of improved signal-to-noise ratios (e.g., The Thirty Meter Telescope, The Giant Magellan Telescope), or a space-based facility dedicated to the study of such time sensitive UV$-$optical$-$NIR transients.  Existing SN~Ia surveys are currently acting toward optimizing a steady flow of discoveries, while other programs have produced a significant number of publicly available spectra \\citep{Richardson01,Matheson08,Blondin12,WISEREP,Silverman12spectra,Folatelli13}. However, for the longterm future we believe it is imperative to begin a discussion of a larger (digital) network of international collaboration by way of (data-) cooperative competition like that done for both The Large Hadron Collider Experiment and Fermi Lab's Tevatron, with multiple competing experiments centered about mutual goals and mutual resources. Otherwise we feel the simultaneous collection of even very high quality temporal datasets by multiple groups will continue to create an inefficient pursuit of over-observing the most high profile event(s) of the year with a less than complete dataset.  Such observational pursuits require an increasingly focused effort toward observing bright and nearby events. For example, 206 supernovae were reported in 1999 and 67 were brighter than 18th magnitude while only three reached $\\sim$ 13 magnitude\\footnote[21]{Quoted from the archives page of \\url{http://www.rochesterastronomy.org/snimages/}.}. By 2012 the number of found supernovae increased to 1045 while 78 were brighter than 16th magnitude and five brighter than 13th magnitude. This clearly indicates that supernovae caught early are more prevalent than $\\sim$ 15 years ago and it is worthwhile for multiple groups to continually increase collaborative efforts for the brightest events. Essentially this could be accomplished without interfering with spectrum-limited high[er]-z surveys by considering a distance threshold ($\\lesssim$ 10$-$30 Mpc) as part of the public domain. Additionally, surveys that corroborate the immediate release of discoveries would further increase the number of well-observed events and could be supplemented and sustained with staggered observations given that there are two celestial hemispheres, unpredictable weather patterns, and caught-early opportunities nearly every week during active surveying.  In conclusion, to extract details of the spectroscopic behavior for all SN~Ia subtypes, during all phases, larger samples of {\\it well-observed} events are essential, beginning from as close to the onset of the explosion as possible (e.g., SN~1999ac, 2009ig, 2011fe, 2012cg, 2012fr), where SN~Ia homogeneity diverges the most (see \\citealt{Zheng13} for the most recent instance in SN 2013dy). Near-continuous temporal observations are most important for at least the first 1$-$2 months post-explosion and biweekly to monthly follow-up thereafter for $\\sim$ 1 year. SN~Ia spectra are far too complicated to do so otherwise. Even normal SN~Ia deserve UV$-$optical$-$IR spectroscopic follow-up at a 1:1 to 2:1 ratio between days passed and spectrum taken, whenever possible, given that fine differences between normal SN~Ia detail the variance in explosion mechanism parameters and initial conditions of their unobserved progenitor systems. It is through such observing campaigns that the true diversity to the underlying nature of SN~Ia events will be better understood."
    },
    "1402/1402.6101_arXiv.txt": {
        "abstract": "With the halo catalog from the {\\it Millennium Simulation}, we analyze the weak-lensing measured density profiles for clusters of galaxies, paying attention to the determination of the concentration--mass ($c$--$M$) relation which can be biased by the center offset, selection effect, and shape noise from intrinsic ellipticities of background galaxies. Several different methods of locating the center of a cluster from weak-lensing effects alone are explored. We find that, for intermediate redshift clusters, the highest peak from our newly proposed two-scale smoothing method applied to the reconstructed convergence field, first with a smoothing scale of $2\\arcmin$ and then $0\\overset{\\prime}{.}5$, corresponds best to the true center. Assuming the parameterized Navarro--Frenk--White profile, we fit the reduced tangential shear signals around different centers identified by different methods. It is shown that, for the ensemble median values, a center offset larger than one scale radius $r_s$ can bias the derived mass and concentration significantly lower than the true values, especially for low-mass halos. However, the existence of noise can compensate for the offset effect and reduce the systematic bias, although the scatter of mass and concentration becomes considerably larger. Statistically, the bias effect of center offset on the $c$--$M$ relation is insignificant if an appropriate center finding method is adopted. On the other hand, noise from intrinsic ellipticities can bias the $c$--$M$ relation derived from a sample of weak-lensing analyzed clusters if a simple $\\chi^2$ fitting method is used. To properly account for the scatter and covariance between $c$ and $M$, we apply a Bayesian method to improve the statistical analysis of the $c$--$M$ relation. It is shown that this new method allows us to derive the $c$--$M$ relation with significantly reduced biases. ",
        "introduction": "In the cold dark matter paradigm, dark matter halos play critical roles in hosting the formation of luminous objects and shaping the observable universe. The structures of dark matter halos themselves also carry important cosmological information closely related to properties of dark matter particles, as well as different astrophysical processes affecting their formation and evolution. Numerical simulations reveal an approximately universal behavior of the density profile with $\\rho \\propto r^{-1}$ at the inner part and $\\rho \\propto r^{-3}$ at the outer part of halos. Different fitting models for the density profile have been proposed, such as Navarro--Frenk--White (NFW; \\citealt{1996ApJ...462..563N,1997ApJ...490..493N}), generalized NFW \\citep{1990ApJ...356..359H,1996MNRAS.278..488Z} and Einasto profiles \\citep{1965TrAlm...5...87E,2012A&A...540A..70R}. Among others, the profile proposed by \\citet{1996ApJ...462..563N,1997ApJ...490..493N}, namely, the NFW profile, has been widely adopted to describe the density distribution. It can be written as \\begin{equation}\\label{eq:rho} \\rho(r)=\\frac{\\rho_s}{(r/r_s)(1+r/r_s)^2}, \\end{equation} where $\\rho_s$ and $r_s$ are the characteristic density and scale of a halo respectively. Given the halo radius \\rdlt, another equivalent set of parameters, the concentration $\\cdlt$ and the mass $\\mdlt$, is normally used to characterize the density profile, where $\\cdlt=\\rdlt/r_s$ and $\\mdlt=(4\\pi/3)\\Delta\\rho_{\\rm crit}\\rdlt^{3}$. Here $\\Delta$ is the average overdensity parameter within $\\rdlt$ with respect to the critical matter density $\\rho_{\\rm crit}$ of the universe, and $\\Delta=200$ or $\\Delta=\\Delta_{\\rm vir}$ for the virialized region of a halo are often adopted in different analyses.  The relation between $\\cdlt$ and $\\mdlt$ ($c$--$M$ relation) is related to the formation history of dark matter halos and has been studied extensively with simulations \\citep[e.g.,][]{2001MNRAS.321..559B,2008MNRAS.390L..64D,2008MNRAS.387..536G,2009ApJ...707..354Z,2012MNRAS.423.3018P,2014MNRAS.441..378L}. At low redshifts, the relation of the two can be approximately described by a power law given by \\begin{equation}\\label{eq:cm} \\lgg\\,\\cdlt=\\lgg\\,A(z)+\\alpha(z) \\lgg\\,\\frac{\\mdlt}{M_p} \\end{equation} in log-space where both $A(z)$ and $|\\alpha(z)|$ are decreasing functions of redshift \\citep{2008MNRAS.387..536G,2011MNRAS.411..584M}. With certain variations for the results obtained from one group to another, simulation studies find the slope $\\alpha\\sim -0.1$ at redshift $z\\sim 0$. A is more sensitive to cosmological models than $\\alpha$ and depends on the definition of the halo mass (such as $\\Delta=200$ or $\\Delta_{\\rm vir}$) and choice of the pivot mass $M_p$ \\citep[e.g.,][]{2004A&A...416..853D,2006ApJ...646..815S,2008MNRAS.390L..64D,2011MNRAS.411..584M,2012MNRAS.424.1244F,2013ApJ...766...32B}.  Because of its cosmological significance, extensive efforts have been made to observationally measure the density profile of dark matter halos and further investigate their $c$--$M$ relation. For that, clusters of galaxies are the most important targets. They are known as the largest virialized objects in the Universe, and their formation and evolution are dominantly affected by gravitational processes. Among others, gravitational lensing analyses provide a direct way to study the dark matter distribution of clusters of galaxies \\citep[e.g.,][]{2001PhR...340..291B,2002MNRAS.337.1269W,2012ApJ...748...56S,2011A&ARv..19...47K, 2013SSRv..177...75H}.  Concerning the $c$--$M$ relation from lensing studies, \\citeauthor{2010PASJ...62..811O} (2010, hereafter Ok10) perform detailed weak-lensing analyses for $30$ X-ray selected clusters and show that the power index of the $c$--$M$ relation is $\\alpha\\sim -0.4$, significantly steeper than the simulation result. \\citeauthor{2012MNRAS.420.3213O} (2012, hereafter Og12) carry combined strong and weak-lensing studies for $28$ strong-lensing selected clusters and obtain a slope of $\\alpha\\sim -0.59$, noting that the expected $\\alpha$ is $\\sim-0.2$  by taking into account the strong-lensing selection bias. Combining the richness measurements and the Einstein radii from strong arcs, \\citet{2012ApJ...761....1W} study $10$ strong-lensing clusters with redshift $0.26\\le z\\le 0.56$. They find $\\alpha\\sim -0.45$ and show that the steep slope is mainly driven by the low-mass clusters in their sample. \\cite{2013MNRAS.434..878S} compile the weak-lensing analyses of $31$ massive clusters at high redshift with $z\\sim (0.8, 1.5)$, and also find a steep $\\alpha\\sim -0.8$. It has been suggested that the steep slope from observations may indicate that physical effects, such as baryon cooling and the dynamical status of clusters, play non-negligible roles in affecting the mass distribution of dark matter halos \\citep[e.g.,][]{2008JCAP...08..006M,2010MNRAS.405.2078M,2010A&A...524A..68E,2012MNRAS.420.3213O,2012MNRAS.424.1244F,2013MNRAS.434..878S}.  On the other hand, \\citet{2007MNRAS.379..190C} combine a variety of observational results with their $10$ strong lensing clusters and find a slope of $\\alpha\\sim -0.14$, consistent with simulation results. Based on data from the Sloan Digital Sky Survey (SDSS; \\citealt{2000AJ....120.1579Y}), \\citet{2008JCAP...08..006M} perform stacked weak-lensing analyses around galaxies binned by their luminosities and MaxBCG clusters binned by their richness. The mass range covers $\\sim 10^{12}\\msun$ to $\\sim 10^{15}\\msun$. They also show that their results are consistent with the $c$--$M$ relation with the slope parameter $\\alpha\\sim -0.13$.  Noticing considerable differences from different analyses, in order to compare with simulation results and draw conclusions properly, it is therefore necessary to understand how various effects can influence the lensing derived density profile and mass, and further the $c$--$M$ relation of dark matter halos. Focusing on weak-lensing studies, in this paper, we specifically investigate how the $c$--$M$ relation is affected by center choice of clusters and shape noise of background galaxies using the halo catalog from the {\\it Millennium Simulation} (MS; \\citealt{2005Natur.435..629S}).  The center identification for a cluster of galaxies is important in weak-lensing analyses. Optical or X-ray observations are commonly used for locating the center defined as the position of the brightest cluster galaxy (BCG) or the X-ray centroid \\citep[e.g.,][]{2007ApJ...660..221K,2007ApJ...660..239K,2010A&A...524A..68E,2012ApJ...757....2G}. Regardless of the very possible offsets between these observables and true halo centers \\citep{2006A&A...451..395C,2012MNRAS.426.2944Z,2012MNRAS.420.2120M}, misidentification of, e.g., BCG, can result in a large fraction of offcenters from true ones \\citep[e.g.,][]{2007arXiv0709.1159J,2012MNRAS.426.2944Z}. On the other hand, for massive clusters with high-quality, weak-lensing observations, it is possible to find the centers of individual clusters self-consistently using lensing data alone \\citep[e.g.,][]{2010MNRAS.405.2215O, 2010A&A...520A..58I, 2012A&A...546A..79I}. Expectedly, tracing the potential center of a cluster is the ideal case; however, in reality, the weak-lensing determined center can be offset from the true center due to various effects \\citep[e.g.,][]{2006A&A...451..395C, 2010ApJ...719.1408F}. With simulated clusters, \\cite{2012MNRAS.419.3547D} demonstrate that the noise from intrinsic ellipticities of source galaxies and the smoothing procedures dominantly lead to the offsets.  The general conclusion is that the center offset can affect the determination of the density profile, particularly on small scales. For example, the halo mass derived from the stacked lensing signals can be systematically low if the center offset effects are not properly considered \\citep[e.g.,][]{2010MNRAS.405.2078M, 2012ApJ...757....2G}. In this paper, we systematically explore four different methods of locating weak-lensing centers, one for shear field and three for convergence maps, using the simulated clusters from MS. Their influences on the weak-lensing-derived $c$--$M$ relation are statistically evaluated .   In addition, we also pay particular attention to the effects of shape noise from intrinsic ellipticities emphasizing their impacts on the weak-lensing-derived $c$--$M$ relation. Significant studies have been done to understand the uncertainties in weak-lensing determination of cluster mass and profile due to different effects, including the nonsphericities of cluster mass distributions and the substructures therein, the projection effects of correlated and uncorrelated large-scale structures, profile fitting methods, the shape noise of source galaxies, etc. \\citep[e.g.,][]{2007MNRAS.380..149C, 2010MNRAS.405.2078M, 2011MNRAS.412.2095H, 2011MNRAS.414.1851O, 2011ApJ...740...25B, 2012MNRAS.421.1073B, 2012MNRAS.426.1558G, 2014MNRAS.440.1899G}. Most of these studies concentrate on the bias and scatters themselves. Some of them discuss the consequent impact on weak-lensing-derived $c$--$M$ relation without systematically analyzing how the impact depends on the level of scatter and their correlation \\citep{2012MNRAS.421.1073B, 2012MNRAS.426.1558G, 2014MNRAS.440.1899G}.  In our study, we perform systematic analyses about the impact of the shape noise on the $c$--$M$ relation. By taking into account the scatter in concentration and mass and covariance between these two variables, we use a Bayesian method that includes selection effect based on halo mass function to infer the realistic $c$--$M$ relation from weak-lensing observed $(c,\\,M)$. To account for the sample variance, we also carry out Monte Carlo analyses with respect to $19$ clusters of Ok10 and $28$ clusters of Og12 and compare with their observational results on the $c$--$M$ relation.  The rest of the paper is organized as follows. We present the mock data generation for weak-lensing analyses in Section 2. Different center identification methods and profile fitting are described in Section 3. Section 4 contains the detailed results. Discussions are given in Section 5. ",
        "conclusions": "Based on the halo catalog from the {\\it Millennium Simulation}, we generate mock weak-lensing data for $\\sim 2000$ clusters and systematically investigate the effects of different center finding methods and the noise on the weak-lensing-derived $c$--$M$ relation.  Four different methods of center identification using weak-lensing data alone are explored. In agreement with other studies \\citep[e.g.,][]{2012MNRAS.419.3547D}, we find that both the shape noise and the smoothing can lead to misidentifications of centers. For high-mass halos, their intrinsically irregular mass distribution can also considerably affect the center identifications. In general, the performance of a center identification method in a weak-lensing regime mainly depends on the ${\\rm S/N}$. For intermediate redshift clusters, the quantity $E_{\\rm SN}=(D_{ds}/D_s)/\\sigma_n$ can be used as an indicator to estimate the efficiency of weak-lensing center identification methods. Among different methods, our newly proposed two-scale smoothing method, 2K05K, performs the best. Our simulation analyses show that for $E_{\\rm SN}\\sim10$, at least $90\\%$ of the identified centers by 2K05K have offsets to the true centers being less than $r_s$, demonstrating its great potential in weak-lensing cluster studies.  For NFW profile constraints, consistent with other studies \\citep[e.g.,][]{2010MNRAS.405.2078M, 2012ApJ...757....2G}, we notice that the center offset alone can lead to a large negative bias to the fitted mass and concentration parameter if it is larger than $r_s$. On the other hand, unlike studies in which centers are defined independent of weak-lensing data, we identify centers from weak-lensing analyses. Therefore, the located centers do have high apparent weak-lensing signals, either associated with real subclumps for massive halos or due to the chance alignments of source galaxies leading to high false peaks for relatively low-mass halos. Consequently, with these contributions, the net biases on the mass and concentration parameter are actually smaller than that from the pure effect of center offset without including these additional signals.  Concentrating on statistically analyzing the $c$--$M$ relation from a sample of weak-lensing studied clusters, each with an NFW fitted $(c,M)$, we find that for our default case with $n_g=30\\hbox{ arcmin}^{-2}$ and $\\sigma_{\\epsilon_s}=0.4$, using centers identified by 2K05K gives rise to almost the same result on the $c$--$M$ relation as that using the true centers in the NFW profile fitting. Again, this shows the feasibility of finding centers in weak-lensing cluster analyses.   Regarding the shape noise effects, our systematic analyses, assuming different noise levels, reveal that they can significantly affect the apparent $c$--$M$ relation from weak-lensing studies. If the simple $\\chi^2$ fitting method is adopted in deriving the $c$--$M$ relation from a sample of weak-lensing analyzed clusters, we generally obtain a steeper slope parameter $\\alpha$, in line with a number of observational studies \\citep[e.g.,][Og12, Ok10]{2013MNRAS.434..878S}. This is mainly due to the scatter associated with the weak-lensing determined $(c,M)$ and their covariance, which are not properly accounted for in the simple $\\chi^2$ fitting.  For the stacking analyses based on the weak-lensing-derived mass, similar problems exist; thus similar noise effects occur and lead to a steeper $c$--$M$ relation than the true one. On the other hand, if the stacking is based on true mass, the derived $c$--$M$ relation shows no significant bias with respect to the true $c$--$M$ relation. Observationally, we do not directly know the true mass of halos. However, if we have other weak-lensing independent observables that strongly correlate with the underlying true mass of halos, they can be used to group clusters for stacking. We then expect no significant bias in the derived $c$--$M$ relation. For that, the galaxy-galaxy lensing technique can be advantageous for either the luminosity or the stellar mass of foreground galaxies, which are often used to form bins for stacking analyses \\citep[e.g.,][]{2006MNRAS.368..715M, 2009MNRAS.394.1016L}. They are shown to be good tracers of their host halo mass \\citep{2007ApJ...671..153Y}. The galaxy-galaxy lensing analyses of \\citet{2008JCAP...08..006M} indeed give rise to a $c$--$M$ relation that is rather consistent with the 3D result with a slope of $\\alpha\\sim -0.1$.  On the other hand, concerning weak-lensing analyses alone, and the apparent steep $c$--$M$ relation seen from the simple $\\chi^2$ fitting, we further develop a Bayesian method in probing the $c$--$M$ relation, taking into account $(c,M)$ scatter and covariance. The sample selection effect is also considered by invoking a lower mass limit for a mass selected sample. In principle, more complicated selection functions suitable for different samples can be incorporated into the Bayesian method. Our results show that with the Bayesian method, we can extract the underlying $c$--$M$ relation of dark matter halos faithfully from the apparently noise-affected data for a wide range of noise levels.  Having finished our analyses, we note a recent study by \\citet{2013MNRAS.436..503A}. They perform constraints on the $c$--$M$ relation by combining the information of Einstein radius and optical richness for $26$ group and cluster-scale objects also using a Bayesian approach. They also point out that the steep slope found by some previous studies is likely due to the degeneracy between concentration and mass.   To account for the sample variance of samples containing a limited number of clusters, we conduct Monte Carlo analyses with respect to the samples of Ok10 and Og12. We find that the noise effects significantly contribute to the apparently steep $\\alpha$ in the $c$--$M$ relation obtained by Ok10 and by Og12. Our Bayesian analyses show that if a lower mass limit $M_{\\rm low}=3.57\\times10^{14}h^{-1}\\msun$ is adopted, the Ok10 result is reasonably consistent with that of the MS clusters within $1-\\sigma$ range. For Og12, while the slope $\\alpha$ is consistent with that of MS clusters including the noise effects, the amplitude parameter $A$ in the $c$--$M$ relation is higher from Og12, even with the strong-lensing bias taken into consideration.  It is noted that MS is a dark matter-only simulation without including baryonic physics. It is believed that baryonic cooling and star formation tend to increase the concentration of the halo mass distribution \\citep[e.g.,][]{2010MNRAS.405.2161D,2012MNRAS.426.1558G,2011MNRAS.416.2539K}. On the other hand, feedback from supernova and AGNs can heat the gas and suppress star formation. Therefore, the overall effects of the baryonic component on the density profile of cluster-scale halos may not be significant \\citep[e.g.,][]{2010MNRAS.405.2161D,2010MNRAS.406..434M,2012MNRAS.424.1244F}.   It should also be pointed out that MS adopts the cosmological model with a higher $\\sigma_8=0.9$ and lower $\\Omega_\\smM=0.25$ than the current best model from latest observations, \\eg {\\it WMAP}7 and {\\it Planck}  \\citep{2013arXiv1303.5076P,2011ApJS..192...18K}. Varying cosmological models can affect the properties of dark matter halos and therefore their $c$--$M$ relation \\citep{2003ApJ...598...49Z,2011MNRAS.418.2422R}. It is shown that while the slope $\\alpha$ is not sensitive to cosmological models, the amplitude $A$ is \\citep[e.g.,][]{2004A&A...416..853D,2008MNRAS.391.1940M,2008MNRAS.390L..64D}. For instance, $A$ for {\\it WMAP}1 cosmology is about $19\\%$ higher than that of {\\it WMAP}5 model, and the slope $\\alpha$ is consistent with $-0.1$ in both models \\citep{2008MNRAS.391.1940M,2008MNRAS.390L..64D}.  However, very recent studies perform a direct comparison of halo concentration from the MS and predict from the {\\it Planck} cosmological model \\citep{2014MNRAS.441..378L}. The results show that for cluster-scale halos, the concentration parameter is nearly the same for the two models. This is because the {\\it Planck} cosmological model has a larger $\\Omega_m$ than that of MS, although $\\sigma_8$ is smaller.   Our analyses for massive halos show that substructures can affect the center identifications, and also lead to a negative bias to 2D concentration in comparison with that of 3-D halos \\citep[see also][]{2012MNRAS.426.1558G,2012MNRAS.421.1073B}. It has been shown that a halo with lower concentration and higher mass and a halo formed later tend to contain a higher fraction of mass in substructures \\citep{2004MNRAS.355..819G,2010MNRAS.404..502G, 2012MNRAS.420.2978C}. The change of cosmology is expected to affect the mass accretion history and therefore the subhalo abundances. However, recent analyses of \\citet{2012MNRAS.425.2169G} and \\citet{2012MNRAS.424.2715W} show that substructure abundances are not very sensitive to cosmological models. For example, \\citet{2012MNRAS.424.2715W} find that subhalo abundances at given host halo mass are very similar between MS and {\\it WMAP}7 cosmology, although MS has a higher $\\sigma_8$.  Overall, our results presented here based on MS-simulated clusters may not be seriously affected by the specific cosmological model adopted by MS. The reasonable agreement between our simulation analyses, including the noise and selection effects and that of Ok10 for the $c$--$M$ relation, may indicate that our current understanding of the structure formation is well on track. On the other hand, the somewhat larger $A$ seen in Og12 still deserves further studies.  Apart from comparing with Ok10 and Og12, in this paper, we emphasize the methodology of center identification using weak-lensing analyses alone, and the noise effects on the weak-lensing-derived $c$--$M$ relation. Future weak-lensing observations can result in a large number of clusters with high signal-to-noise ratios. We expect that both our newly proposed 2K05K center identification method, and the Bayesian method in deriving the $c$--$M$ relation from a sample of weak-lensing studied clusters can have important applications in future cosmological studies."
    },
    "1402/1402.2980_arXiv.txt": {
        "abstract": "We present results of our time variability studies of \\mgii\\ and \\aliii\\ absorption lines in a sample of 22 Low Ionization Broad Absorption Line QSOs (LoBAL QSOs) at 0.2 $\\leq$ z$_{em}$ $\\leq$ 2.1  using the 2m telescope at IUCAA Girawali Observatory  over a time-scale of 10 days to 7.69 years in the QSO's rest frame.  Spectra are analysed in conjunction with photometric light curves from Catalina Real-Time Transient Survey. Long time-scale (i.e $\\ge$ 1 year) absorption line variability is seen in 8 cases (36\\% systems) while only 4 of them (i.e 18\\% systems) show variability over short time-scales (i.e  $<$ 1 year). We notice a tendency of highly variable  LoBAL QSOs to have high ejection velocity, low equivalent width and low redshift. The detection rate of variability in  LoBAL QSOs  showing Fe fine-structure lines  (FeLoBAL QSOs) is less than that seen in non-Fe  LoBAL QSOs. Absorption line variability is more frequently detected in QSOs having continuum dominated by Fe emission lines compared to rest of the QSOs. Confirming these trends with a bigger sample will give vital clues for understanding the physical distinction between different BAL QSO sub-classes. We correlate the absorption line variability with various parameters derived from continuum light curves and  find no clear correlation between continuum flux and absorption line variabilities. However, sources with large absorption line variability also show large variability in their light curves. We also see appearance/disappearance of absorption components  in 2 cases and clear indications for profile variations in 4 cases.  The observed variability can be best explained by a combination of process driven by continuum variations and clouds transiting across the line of sight. ",
        "introduction": "In recent years, the importance of understanding the nature and origin of QSO outflows  has become widely recognized in the context of  well studied  correlations between the Super Massive Black Hole (SMBH) and host galaxy properties \\citep{silk98,ferrase00}. The inferred kinetic luminosities of the these outflows suggest that they could in principle be a major contributor to the  Active Galactic Nuclei (AGN) feedback mechanisms \\citep{moe09,dunn10,bautista10,aoki11,borguet12,borguet13}. QSO outflows are also an important diagnostic of the physical conditions and chemical enrichment prevailing in the inner regions of the AGNs.  QSO outflows are readily identified as  blue shifted broad absorption lines (BAL) seen in 10-40\\%  of the total quasars \\citep{hewett03,reichard03,trump06,allen11,paris12}. By definition, BAL QSOs are characterized by  absorption troughs which dip by more than 10\\% below the continuum for a contiguous stretch of at least 2000 \\kms and extend up to velocities $\\sim$ 0.2 c with respect to the systemic redshift \\citep{weymann91}. BAL QSOs are further divided in three sub-types depending on the type of transition seen in absorption. Most of the known BAL QSOs are high ionization BALs (HiBALs) characterized by absorptions from \\civ, \\siiv, \\nv, and \\lya. About 15\\% are Low ionization BALs  (LoBAL QSOs) which exhibit absorption lines of \\mgii, \\alii, \\aliii\\ and sometimes \\feii\\ and \\feiii\\ in addition to the above listed high ionization absorption lines. The rarer ($\\sim$ 1\\%) iron  LoBAL QSOs (FeLoBAL QSOs), also show absorption from excited fine-structure levels of \\feii\\ or \\feiii. Recently thanks to the cosmic origin spectrograph on-board HST (COS/HST), outflows traced by very highly ionized species like \\nevii, \\mgx\\ and \\nax\\ have been detected in $\\sim$ 40\\% of low-z QSOs \\citep{muzahid12,muzahid13}.  The observed  BAL incidence is attributed to either an orientational effect where the observer's line of sight passes through the BAL clouds or an evolutionary phase in the QSO lifetime or a combination of the two.  Even though recent spectroscopic surveys have dramatically increased the number of known BAL QSOs, still there is no commonly accepted  model which explains the geometry and the physical properties of the outflows. Line driven radiative  acceleration is suggested based on spectroscopic signatures \\citep{arav94,arav95,anand02}. \\citet{dekool95} proposed a model in which gas is clumped in small clouds confined by magnetic pressure and driven by radiation-pressure.  There is another dynamical model which assumes outflows as winds streaming out of the accretion-disk at a distance R $\\sim$ 10$^{16}$ cm and driven by the radiative UV line pressure \\citep{murray95,proga00,elvis00}. The proposed   \"shielding gas\"  absorbs the X-rays from the quasar and protects the wind from being completely ionized, also explains the observed X -ray weakness of BAL QSOs \\citep[see,][]{green95,gallagher06,fan09,gibson09,stalin11}. Most current models favor the equatorial geometry of the BAL outflows although after the discovery of radio loud BAL QSOs \\citep{becker97,becker00}  polar outflows have been suggested for some of them \\citep{ghosh07}. A two-component wind model  combining the polar and equatorial components have also been suggested to explain the polar outflows \\citep{lamy04,proga04,borguet10}. The exact geometry of the BAL outflows is still under theoretical and observational investigation. \\begin{table*} \\caption{Log of observations} \\flushbottom \\begin{tabular}{|c|c|c|c|c|c|c|c|} \\hline QSO&Instrument& Date&MJD&Exposure Time&$\\lambda$ Coverage&Resolution&S/N$^a$\\\\ &&(DD-MM-YYYY)&&(mins)&($\\AA$)&(kms$^{-1}$)&\\\\ \\hline & SDSS        & 23-09-2000 & 51812 & 15x6 & 3800-9200 & 150 & 34\\\\ & SDSS        & 20-11-2000 & 51868 & 15x5 & 3800-9200 & 150 & 34 \\\\ & SDSS        & 01-12-2000 & 51879 & 15x4 & 3800-9200 & 150 & 34\\\\ J004527.68+143816.1 & IGO/IFOSC 7 & 13-12-2007 & 54447 & 40x2 &\t3800-6840 & 310 & 11\\\\ & IGO/IFOSC 7 & 18-12-2008 & 54818 & 40x2 &\t3800-6840 & 310 & 19\\\\ & IGO/IFORS 1 & 22-12-2009 & 55187 & 35x3 & 3270-6160 & 370 & 12 \\\\ & IGO/IFOSC 8 & 01-01-2009 & 54832 & 45x3 & 5800-8350 & 240 & 17  \\\\ \\hline & SDSS        & 06-09-2000 & 51793 & 15x3 & 3800-9200 & 150 & 30\t\\\\ J014950.96-010314.1 & IGO/IFOSC 7 & 09-01-2008 & 54474 & 45x2 & 3800-6840 & 310 & 11      \\\\ & IGO/IFOSC 7 & 23-12-2008 & 54823 & 45x5 & 3800-6840 & 310 & 12      \\\\ & SDSS        & 26-08-2010 & 55447 & 15x5 & 3800-9200 & 150 & 23\\\\ & IGO/IFOSC 7 & 13-02-2012 & 55970 & 45x3 & 3800-6840 & 310 & 11      \\\\ \\hline & SDSS        & 29-09-2000 & 51816 & 45x1 & 3800-9200 & 150 & 44    \\\\ J030000.57+004828.0 & SDSS        & 29-11-2000 & 51877 & 50x1 & 3800-9200 & 150 & 40    \\\\ (FeLoBAL QSO)\t    & SDSS        & 23-10-2001 & 52205 & 45x1 & 3800-9200 & 150 & 33     \\\\ & IGO/IFOSC 7 & 12-12-2007 & 54446 & 35x2 & 3800-6840 & 310 &  5     \\\\ & IGO/IFOSC 7 & 27-12-2008 & 54817 & 45x8 & 3800-6840 & 310 & 54      \\\\ \\hline & SDSS        & 15-01-2001 & 51924 & 80x1 & 3800-9200 & 150 & 27 \\\\ & IGO/IFOSC 8 & 14-12-2007 & 54448 & 45x3 & 5800-8350 & 240 & 10  \\\\ & IGO/IFOSC 7 & 09-01-2008 & 54474 & 45x3 & 3800-6840 & 310 & 23   \\\\ J031856.62-060037.7 & IGO/IFOSC 8 & 27-11-2008 & 54797 & 45x3 & 5800-8350 & 240 & 21    \\\\ (FeLoBAL QSO)       & IGO/IFOSC 7 & 31-01-2009 & 54862 & 45x4 & 3800-6840 & 310 & 19     \\\\ & IGO/IFOSC 8 & 23-12-2009 & 55188 & 45x4 & 5800-8350 & 240 & 19    \\\\ & IGO/IFOSC 7 & 24-12-2009 & 55189 & 45x4 & 3800-6840 & 310 & 17    \\\\ \\hline & SDSS        & 01-01-2001 & 51910 & 15x3 &\t3800-9200 & 150 & 33\t\\\\ J033438.28-071149.0 & IGO/IFORS 1 & 20-01-2010 & 55216 & 45x4 & 3270-6160 & 370 & 28 \\\\ & IGO/IFORS 1 & 24-01-2012 & 55950 & 45x2 & 3270-6160 & 370 &  7\\\\ \\hline & SDSS\t  & 29-11-2000 & 51877 & 15x3 & 3800-9200 & 150 & 37\t\\\\ & IGO/IFOSC 7 & 21-12-2006 & 54090 & 30x2 & 3800-6840 & 310 & 6     \\\\ & IGO/IFOSC 7 & 04-12-2008 & 54804 & 45x3 & 3800-6840 & 310 & 29     \\\\ J073739.96+384413.2 & IGO/IFORS 1 & 24-12-2009 & 55189 & 45x3 & 3270-6160 & 370 & 34     \\\\ & IGO/IFOSC 8 & 05-12-2008 & 54805 & 45x4 & 5800-8350 & 240 & 32    \\\\ & IGO/IFOSC 8 & 31-01-2009 & 54862 & 45x3 & 5800-8350 & 240 & 34    \\\\ & IGO/IFOSC 8 & 25-02-2009 & 54887 & 45x3 & 5800-8350 & 240 & 28    \\\\ \\hline & SDSS        & 24-12-2001 & 52207 & 20x7 & 3800-9200 & 150 & 30      \\\\ J082319.65+433433.7 & SDSS        & 19-02-2001 & 51959 & 15x3 & 3800-9200 & 150 & 21      \\\\ & IGO/IFORS 1 & 13-12-2007 & 54447 & 40x3 & 3270-6160 & 370 &  7      \\\\ & IGO/IFOSC 7 & 06-12-2008 & 54806 & 45x3 & 3800-6840 & 310 & 20    \\\\ & IGO/IFOSC 7 & 19-04-2009 & 54575 & 45x2 & 3800-6840 & 310 & 2       \\\\ & IGO/IFOSC 7 & 25-01-2010 & 55221 & 45x3 & 3800-6840 & 310 & 8    \\\\  \\hline & SDSS        & 19-11-2001 & 52232 & 48x1 & 3800-9200 & 150 & 26   \\\\ J083522.77+424258.3 & IGO/IFORS 1 & 14-12-2007 & 54448 & 45x3 & 3270-6160 & 370 & 9     \\\\ (FeLoBAL QSO)      & IGO/IFOSC 7 & 05-12-2008 & 54805 & 45x3 & 3800-6840 & 310 & 21   \\\\ & IGO/IFOSC 7 & 21-01-2010 & 55217 & 45x3 & 3800-6840 & 310 & 29   \\\\ \\hline & SDSS        & 15-02-2002 & 52320 & 50x1 & 3800-9200 & 150 & 40   \\\\ & IGO/IFOSC 7 & 20-12-2006 & 54089 & 45x3 & 3800-6840 & 310 & 8    \\\\ J084044.41+363327.8 & IGO/IFOSC 7 & 12-12-2007 & 54446 & 45x3 & 3800-6840 & 310 & 23    \\\\ (FeLoBAL QSO)        & IGO/IFOSC 7 & 07-12-2008 & 54807 & 45x4 & 3800-6840 & 310 & 26    \\\\ & IGO/IFOSC 7 & 17-12-2009 & 55182 & 45x6 & 3800-6840 & 310 & 50    \\\\ \\hline & SDSS        & 27-11-2002 & 52605 & 15x6 & 3800-9200 & 150 & 34    \\\\ J085053.12+445122.4 & IGO/IFORS 1 & 19-01-2010 & 55215 & 45x5 & 3270-6160 & 370 & 31     \\\\ & IGO/IFORS 1 & 12-03-2011 & 55632 & 45x5 & 3270-6160 & 370 & 16     \\\\ \\hline & SDSS        & 12-03-2003 & 52710 & 15x3 &\t3800-9200 & 150\t& 45    \\\\ J094443.13+062507.4 & IGO/IFOSC 7 & 03-04-2008 & 54559 & 40x3 & 3800-6840 & 310 & 45    \\\\ & IGO/IFOSC 7 & 30-01-2009 & 54861 & 45x5 & 3800-6840 & 310 & 60    \\\\ \\hline \\end{tabular} \\begin{flushright} \\it{ Continued on next page} \\\\ \\end{flushright} \\label{log} \\end{table*} \\setcounter{table}{0} \\begin{table*} \\caption{{ Continued:} Log of observations} \\flushbottom \\begin{tabular}{|c|c|c|c|c|c|c|c|} \\hline QSO&Instrument& Date&MJD&Exposure Time&$\\lambda$ Coverage&Resolution&S/N$^a$\\\\ &&(DD-MM-YYYY)&&(mins)&($\\AA$)&(kms$^{-1}$)&\\\\ \\hline \\hline & IGO/IFOSC 7 & 24-01-2010 & 55220 & 45x2 & 3800-6840 & 310 & 28    \\\\ J094443.13+062507.4 & IGO/IFOSC 7 & 11-04-2010 & 55297 & 45x2 & 3800-6840 & 310 & 20    \\\\ & IGO/IFOSC 7 & 14-03-2011 & 55634 & 45x4 & 3800-6840 & 310 & 50    \\\\ \\hline & SDSS        & 30-12-2000 & 51908 & 15x3 &\t3800-9200 & 150\t& 27    \\\\ & IGO/IFOSC 7 & 03-04-2008 & 54559 & 40x3 & 3800-6840 & 310 & 22    \\\\ J095232.21+025728.3 & IGO/IFOSC 7 & 25-02-2009 & 54887 & 45x3 & 3800-6840 & 310 & 25    \\\\ & IGO/IFOSC 7 & 25-01-2010 & 55221 & 45x3 & 3800-6840 & 310 & 10    \\\\ & IGO/IFOSC 7 & 15-03-2011 & 55635 & 45x5 & 3800-6840 & 310 & 19    \\\\ \\hline & SDSS        & 12-04-2002 & 52376 & 15x3 &\t3800-9200 & 150\t& 23    \\\\ J101053.98+451817.0 & IGO/IFORS 1 & 09-01-2008 & 54474 & 45x3 & 3270-6160 & 370 & 14     \\\\ & IGO/IFORS 1 & 26-02-2009 & 54888 & 45x3 & 3270-6160 & 370 & 17     \\\\ & IGO/IFORS 1 & 24-01-2010 & 55220 & 45x4 & 3270-6160 & 370 & 18     \\\\ \\hline & SDSS        & 03-01-2003 & 52642 & 15x3 &\t3800-9200 & 150\t& 24   \\\\ J112822.42+482310.0 & IGO/IFOSC 7 & 22-01-2010 & 55218 & 45x3 & 3800-6840 & 310 & 23   \\\\ & IGO/IFOSC 7 & 14-03-2011 & 55634 & 45x4 & 3800-6840 & 310 & 22   \\\\ \\hline & SDSS        & 04-04-2002 & 52368 & 15x3 &\t3800-9200 & 150\t& 25    \\\\ J114340.96+520303.3 & IGO/IFORS 1 & 08-01-2008 & 54473 & 40x2 & 3270-6160 & 370 & 11     \\\\ & IGO/IFORS 1 & 28-01-2009 & 54859 & 45x2 & 3270-6160 & 370 & 13     \\\\ \\hline & SDSS        & 25-04-2001 & 52024 & 15x6 &\t3800-9200 & 150\t& 62    \\\\ J120813.42+023015.1 & IGO/IFOSC 7 & 23-01-2010 & 55219 & 40x3 & 3800-6840 & 310 & 57  \\\\ & IGO/IFOSC 7 & 13-03-2011 & 55633 & 45x3 & 3800-6840 & 310 & 42   \\\\ \\hline & SDSS        & 28-04-2000 & 51662 & 15x3 & 3800-9200 & 150 & 27\\\\ & SDSS        & 15-02-2001 & 51955 & 15x7 &\t3800-9200 & 150\t& 40     \\\\ & IGO/IFOSC 7 & 03-04-2008 & 54559 & 40x2 & 3800-6840 & 300 & 11\\\\ J133356.02+001229.1 & IGO/IFORS 1 & 26-02-2009 & 54888 & 45x3 & 3270-6160 & 360 & 21\\\\ & IGO/IFOSC 7 & 26-03-2009 & 54916 & 45x5 & 3800-6840 & 300 & 15\\\\ & IGO/IFORS 1 & 22-01-2010 & 55218 & 45x5 & 3270-6160 & 360 & 26\\\\ & IGO/IFOSC 7 & 06-04-2011 & 55657 & 45x3 & 3800-6840 & 300 & 19\\\\ \\hline & SDSS        & 01-06-2002 & 52426 & 15x3 & 3800-9200 & 150 & 30  \\\\ J133428.06-012349.0 & IGO/IFOSC 7 & 04-04-2008 & 54560 & 40x3 & 3800-6840 & 310 & 18  \\\\ & IGO/IFOSC 7 & 23-12-2008 & 54823 & 45x5 & 3800-6840 & 310 & 30   \\\\ \\hline & SDSS        & 21-04-2001 & 52026 & 15x8 &\t3800-9200 & 150\t& 53    \\\\ J144842.45+042403.1 & IGO/IFORS 1 & 01-04-2010 & 55287 & 45x2 & 3270-6160 & 370 & 9     \\\\ & IGO/IFOSC 8 & 22-03-2010 & 55277 & 45x2 & 5800-8350 & 240 & 12    \\\\ & IGO/IFOSC 7 & 15-03-2011 & 55635 & 45x3 & 3800-6840 & 310 & 22  \\\\ \\hline & SDSS        & 05-05-2003 & 52764 & 15x3 &\t3800-9200 & 150\t& 33    \\\\ J161425.17+375210.7 & IGO/IFOSC 7 & 03-04-2008 & 54559 & 40x1 & 3800-6840 & 310 & 19    \\\\ & IGO/IFOSC 7 & 11-04-2010 & 55297 & 45x3 & 3800-6840 & 310 & 25    \\\\ & IGO/IFOSC 7 & 07-04-2011 & 55658 & 45x1 & 3800-6840 & 310 & 27    \\\\ \\hline  & SDSS        & 04-09-2000 & 51804 & 53x1 & 3800-9200 & 150 & 41  \\\\ & VLT/FORS1   & 20-09-2003 & 52902 &  5x1 & 3300-11000& 370 & 62  \\\\ J221511.93-004550.0 & IGO/IFOSC 7 & 31-10-2008 & 54783 & 45x4 & 3800-6840 & 310 & 38  \\\\ (FeLoBAL QSO)        & IGO/IFOSC 8 & 01-11-2008 & 54784 & 45x4 & 5800-8350 & 240 & 39  \\\\ & IGO/IFORS 1 & 04-01-2010 & 55213 & 45x8 & 3270-6160 & 370 & 9   \\\\ & MagE        & 10-08-2010 & 55431 & 15x1 & 3100-10000&  70 & 16  \\\\ \\hline & SDSS        & 12-10-2002 & 52559 & 15x3 &\t3800-9200 & 150\t& 28    \\\\ & SDSS        & 19-08-2007 & 54331 & 15x5 &\t3800-9200 & 150\t& 38    \\\\ J234711.44-103742.4 & IGO/IFOSC 7 & 31-10-2008 & 54770 & 45x3 & 3800-6840 & 310 & 19   \\\\ & IGO/IFOSC 8 & 21-11-2008 & 54791 & 45x4 & 5800-8350 & 240 & 19    \\\\ & IGO/IFOSC 7 & 15-12-2009 & 55180 & 35x2 & 3800-6840 & 310 &  8  \\\\  \\hline \\hline \\end{tabular} \\begin{flushleft} \\it{ $^a$ The values  quoted are mean of S/N calculated per pixel over the wavelength range 5800 \\AA - 6200 \\AA.} \\end{flushleft} \\label{log} \\end{table*}  As the launching radii of the winds in some models are very close to the QSO central engine, instabilities in the accretion disk can result in variations in the outflow properties on time-scales of a few years. Thus, BAL variability studies are important to understand the location and physical conditions in the absorbing gas and the physical mechanisms responsible for these outflows.  BAL troughs are, by and large, not observed to be varying in velocity  at a significant level \\citep{wampler95,vilk01,rupke02}. However, there are several reported cases  of BAL troughs varying in strength: Q1303+308 \\citep{foltz87}, Q1413+113 \\citep{turnshek88}, Q1246-057 \\citep{smith88}, UM 232 \\citep{barlow89}, QSO CSO 203 \\citep{barlow92}, Tol 1037-270 \\citep{anand01}, J1054+0348 \\citep{hamann08} FBQS J1408+3054 \\citep{hall11} and SDSS 1333+0012\\citep{vivek12}.  Apart from these individual reports, there are six studies of BAL variability which looked for variations in a BAL QSO sample. \\citet{barlow94} studied a sample of 23 BAL QSOs covering time-scales of $\\Delta$t$_{rest}$$\\leq$ 1 yr, reported optical depth variability in 15 of them and proposed  the variable photo-ionizing continuum as the primary cause for these variations.  \\citet{lundgren07} analyzed a sample of 29 BAL QSOs with two-epoch  observations in Sloan Digital Sky Survey (SDSS) with similar time-scales and reported similar variabilities.  \\citet{gibson08} studied a sample of 13 BAL QSOs over a longer time-scale ( $\\Delta$t$_{rest}$$\\sim$ 3-6 yrs). Their study found no evidence  for photoionization driven variability and  indicated variations produced by changes in outflow geometry. \\citet{gibson10}, using a sample of  9 BAL QSOs with multi-epoch data, reported that the scatter in the change of absorption equivalent width, $\\Delta$W, increases with the elapsed time between two epochs of observations. \\citet{capellupo11} studied the variability trends in the short-term ($\\le$1~yr in the QSO rest frame) and long-term ($\\ge$ 1 yr in the QSO rest frame) data  for a sample of 24 BAL QSOs and reported variations in 39\\% and 65\\% of the quasars respectively. \\citet{capellupo12} compared the variabilities in the \\civ\\ and \\siiv\\ lines for the above sample and found that \\siiv\\ BALs are more likely to vary than \\civ\\ BALs. \\citet{filiz12} studied the multi-epoch observations of 582 quasars in SDSS and reported 21 examples of BAL trough disappearance in 19 quasars. Very recently, \\citet{filiz13} studied the variability of 428 \\civ\\ and 235 \\siiv\\ BAL troughs identified in multi-epoch SDSS observations and reported a variability fraction of  60\\%. In all these studies, BALs are found not to vary monotonically over time.  It is  also found that  components at higher outflow velocities are more likely to vary than those at lower velocities, and weaker BALs are more likely to vary than stronger BALs. A co-ordinated variability between absorption lines at different velocities are interpreted as  due to changing ionizing radiations whereas variability in  limited portions of broad troughs are identified as movement of individual clouds across our line-of-sight.  Most of the studies favor the movement of clouds across the continuum source over  the changes in the ionizing continuum. The actual scenario may be a complex combination of both.  Till date, all the BAL variability studies  have focused  mainly on variability in \\civ\\ and \\siiv\\ lines. In this paper, we study the nature of time variability in a sample of 22  LoBAL QSOs over the redshift range z$_{em}$$\\sim$ 0.3-2.1 covering time-scales from   10 days to 7.69 years in the quasar rest frame. This work is the third paper in a series on  LoBAL QSO variability. The first paper, \\citet{vivek12} reported the dynamical evolution of the Mg II BAL in SDSS 1333+0012. The second paper,\\citet{vivek12a} studied the time variability of five  FeLoBAL QSOs and reported for the first time, significant variability in the strength of fine structure lines in the source SDSS 2215-0045. Our sample is drawn from the catalog of BAL QSOs from the SDSS 3$^{rd}$  data release \\citep{trump06}. Typically each source, has 3-5 epoch spectroscopic data. Unlike all previous studies, we supplement the spectroscopic data with long-term photometric light curves. This enables us to discuss the connection between the absorption line variability with the continuum flux variability of the QSOs. \\newcommand{\\ag}{$^{\\dagger}$} \\newcommand{\\ap}{$^{\\ddag}$}  \\begin{table*} \\caption{ Source parameters measured from the data.   } \\flushbottom \\begin{tabular}{|l|c|c|c|c|c|c|c|c|c|c|c|} \\hline Name\t\t    & z$_{em}$&\tz$_{abs}$&   V$_{max}$ &  $<w>$ &   L$_{Bol}$&    M$_{BH}$&   $\\Delta$m &  $\\sigma$($\\Delta$m)&slope($\\Delta$m/$\\Delta$t)& Abs. line(s)\\\\ &\t      &          &    (\\kms)   & (\\AA)  & (ergs s$^{-1}$)& (M$_{\\odot}$)&(mag)&(mag)         &               ( mag/year) & probed                         \\\\ (1)\t\t    \t& (2)\t& (3)  & (4)\t  &  (5)\t & (6)\t\t\t  & (7)      & (8)     &  (9)\t&(10)\t   & (11)  \\\\ \\hline \\hline J0045$+$1438    \t&1.992 & 1.948 & 5746.5   &   6.93\t & 3.9$\\times$ 10$^{47}$   & 3.2E+09&  -0.04 &   0.04 & -0.032 $\\pm$  0.007 & \\aliii   \\\\ J0149$-$0103\\ap   \t&1.074 & 1.030 & 7910.7   &   5.60       & 8.4$\\times$ 10$^{46}$   & 6.8E+08&  -0.01 &   0.13 & -0.046 $\\pm$  0.007 & \\mgii   \\\\ J0300$+$0048\\ag\t\t&0.568 & 0.864 & 6492.8   &     -        & 3.8$\\times$ 10$^{45}$   & 3.0E+07&   0.01 &   0.04 & -0.003 $\\pm$  0.024 & \\mgii    \\\\ J0318$-$0600\\ag\t\t&1.968 & 1.941 & 7612.2   &  14.57       & 7.3$\\times$ 10$^{46}$   & 5.9E+08&   0.07 &   0.05 &  0.075 $\\pm$  0.021 & \\mgii,\\aliii    \\\\ J0334$-$0711 \t\t&0.634 & 0.600 & 7537.6   &   6.50       & 4.3$\\times$ 10$^{46}$   & 3.5E+08&     -  &    -   &          -          & \\mgii  \\\\ J0737$+$3844\\ap\t  \t&1.400 & 1.374 & 5000.2   &   7.16       & 1.8$\\times$ 10$^{47}$   & 1.5E+09&  -0.07 &   0.06 & -0.068 $\\pm$  0.004 & \\mgii,\\aliii\\\\ J0823$+$4334 \t\t&1.660 & 1.620 & 11045.0  &   5.44$^a$   & 2.2$\\times$ 10$^{47}$   & 1.8E+09&  -0.03 &   0.04 & -0.036 $\\pm$  0.009 & \\aliii    \\\\ J0835$+$4242\\ag\t\t&0.810 & 0.769 & 1865.9   &   4.96       & 3.9$\\times$ 10$^{46}$   & 3.2E+08&  -0.04 &   0.07 & -0.005 $\\pm$  0.008 & \\mgii    \\\\ J0840$+$3633\\ag\t\t&1.230 & 1.225 & 3880.8   &  32.13       & 8.1$\\times$ 10$^{45}$   & 6.6E+07&   0.03 &   0.04 &  0.015 $\\pm$  0.008 & \\mgii,\\aliii   \\\\ J0850$+$4451 \t\t&0.541 & 0.524 & 5224.1   &  17.22       & 2.1$\\times$ 10$^{46}$   & 1.7E+08&   0.01 &   0.05 &  0.014 $\\pm$  0.006 & \\mgii    \\\\ J0944$+$0625\\ap\t   \t&0.695 & 0.653 & 8507.7   &  15.55       & 1.3$\\times$ 10$^{47}$   & 1.0E+09&   0.04 &   0.04 &  0.010 $\\pm$  0.004 & \\mgii    \\\\ J0952$+$0257 \t\t&1.355 & 1.301 & 18209.2  &  16.99       & 1.5$\\times$ 10$^{47}$   & 1.2E+09&   0.03 &   0.06 & -0.002 $\\pm$  0.007 & \\mgii,\\aliii   \\\\ J1010$+$4518\\ap   \t&1.776 & 1.705 & 12910.7  &   7.21$^a$   & 1.6$\\times$ 10$^{47}$   & 1.3E+09&  -0.02 &   0.07 & -0.014 $\\pm$  0.012 & \\aliii    \\\\ J1128$+$4823 \t\t&0.543 & 0.536 & 5895.8   &  11.82       & 1.5$\\times$ 10$^{46}$   & 1.2E+08&  -0.01 &   0.06 &  0.010 $\\pm$  0.009 & \\mgii    \\\\ J1143$+$5203\\ap\t   \t&1.816 & 1.788 & 18209.2  &  11.64$^a$   & 2.9$\\times$ 10$^{47}$   & 2.3E+09&   0.06 &   0.11 &  0.041 $\\pm$  0.016 & \\aliii    \\\\ J1208$+$0230\\ap   \t&1.180 & 1.067 & 17836.1  &   4.37       & 3.8$\\times$ 10$^{47}$   & 3.1E+09&  -0.04 &   0.03 & -0.032 $\\pm$  0.004 & \\mgii    \\\\ J1333$+$0012 \t\t&0.918 & 0.811 & 20746.5  &   7.04       & 9.4$\\times$ 10$^{46}$   & 7.6E+08&   0.06 &   0.08 &  0.008 $\\pm$  0.005 & \\mgii    \\\\ J1334$-$0123\\ap   \t&1.876 & 1.833 & 18507.7  &  23.66$^a$   & 2.6$\\times$ 10$^{47}$   & 2.1E+09&   0.07 &   0.06 &  0.058 $\\pm$  0.012 & \\mgii,\\aliii    \\\\ J1448$+$0424 \t\t&1.546 & 1.516 & 5821.1   &   8.02       & 3.8$\\times$ 10$^{47}$   & 3.1E+09&  -0.01 &   0.03 & -0.012 $\\pm$  0.010 & \\mgii,\\aliii    \\\\ J1614$+$3752 \t\t&0.553 & 0.525 & 6492.8   &   3.43       & 3.8$\\times$ 10$^{46}$   & 3.1E+08&  -0.04 &   0.08 & -0.026 $\\pm$  0.004 & \\mgii    \\\\ J2215$-$0045\\ag \\ap\t&1.478 & 1.360 & 18582.3  &  21.76       & 2.8$\\times$ 10$^{47}$   & 2.3E+09&  -0.01 &   0.19 & -0.123 $\\pm$  0.009 & \\mgii,\\aliii    \\\\ J2347$-$1037    \t&1.800 & 1.711 & 15522.6  &  27.46       & 2.6$\\times$ 10$^{47}$   & 2.1E+09&  -0.07 &   0.06 & -0.072 $\\pm$  0.008 & \\aliii    \\\\ \\hline \\hline \\end{tabular} \\begin{flushleft} { Notes: Col. 1: QSO name, $^{\\dagger}$ FeLoBAL QSO sources, $^{\\ddag}$ Sources with strong Fe emission.; Col. 2: z$_{em}$ is obtained from the fitting of SDSS composite.; Col. 3: z$_{abs}$ corresponds to the maximum optical depth.; Col. 4: V$_{max}$ is calculated for the \\mgii\\ line from the normalized SDSS spectra. V$_{max}$ is identified as the maximum velocity at which source flux matches with the continuum.; Col. 5: Average \\mgii\\ equivalent widths, $^a$ corresponds to average \\aliii\\  equivalent widths.; Col. 6: L$_{Bol}$ is computed using the prescription, L$_{Bol}$ = 7.9 $\\times \\nu$F$_{\\nu B }$,  of \\citet{marconi04}.;  Col. 7: M$_{BH}$ is derived from L$_{Bol}$ assuming Eddington accretion.; Col. 8:  median $\\Delta$m  values.; Col. 9: standard deviation of $\\Delta$m values.; Col. 10: slope of $\\Delta$m-$\\Delta$t graph.; Col. 11: absorption line(s) probed in this study.}\\\\ \\end{flushleft} \\label{source_param} \\end{table*}   The manuscript is arranged as follows. In Section 2, we provide details of our sample, observations and data reduction. Section 3 describes the BAL variability measurements and Section 4 gives the summary of variability in individual sources.  In Section 5, we explore statistical correlations between absorption line variability and other measurable parameters from spectroscopy and photometric light curves. Discussions and summary of our results are presented in Section 6. Notes on individual variable BAL sources, Light curves  and spectra of all the objects in our data are presented in the Appendix. \\begin{figure*} \\centering \\psfig{figure=plot_delta_revIIp.ps,width=1.0\\linewidth,height=1.2\\linewidth,angle=0}% \\caption{ Results of long-term photometric monitoring of sources in our sample from CRTS. The x-axis represents rest-frame pairwise time lags  and the y-axis represents pairwise magnitude differences for each source measured over different time lags.  The  red/grey line marks the least square fit to the data. The number in  parenthesis indicate the number of independent photometric measurements available for the source. { Blue dotted horizontal lines correspond to magnitude variations of 0.1 magnitude. We consider those objects as variable for which $\\sigma$($\\Delta$m) $>$ 0.1 magnitudes. Variations above this value are seen in the case  of  SDSS J0149-0103 and SDSS J2215-0045. In the case of SDSS J1143+5203, large photometric variations are seen albeit large photometric errors.}} \\label{lc} \\end{figure*} ",
        "conclusions": "We have presented a detailed analysis of a sample of 22  LoBAL QSOs that are monitored both spectroscopically and photometrically  over a period of few days to few years in the QSO rest frame.  Most of the results are based on  equivalent widths, but we show that they are consistent with results from a spectral ratio analysis that avoids problems related to the continuum fitting.  The main results from the study are summarized below. \\begin{enumerate} \\item  Out of 22 BALs monitored, 8 are showing absorption line variability. We define a variable BAL as the one with variations in rest equivalent width $>$ 3$\\sigma$ level and ratio $>$ 5$\\sigma$. \\citet{gibson08} and \\citet{capellupo11}, in their studies on \\civ\\ BALs, reported the fraction of varying BALs to be 92\\% and 67\\% respectively. Very recently, \\citep{filiz13} have found 50-60\\% of \\civ\\ BALs showing variability over a time-scale of 1 - 3.7 years using an extended sample from SDSS.  As compared to the \\civ\\ BALs, the fraction of varying BALs in our overall sample is small (i.e., 36\\%).  \\citet{capellupo11} used a slightly higher threshold (4$\\sigma$) to define the variability and their sample  span the  redshift range 1.2$<$z$<$2.9. In this redshift range, the variability detection is even less (i.e., 25\\%) for our  LoBAL QSO sample. \\item  We  find a greater incidence of variability at long time-scales ($>$ 1 year) than short time-scales ($<$ 1 year). \\citet{capellupo11} find that 39 \\% of \\civ\\ BALs varied in the short-term data, whereas 65 \\% varied in the long-term data.  In this study of \\mgii\\ BALs, we find the respective variable fraction to be 18\\% and 36\\% for short-term and long-term data. This could mean that the physical process causing the BAL variations dominates over multi-year time-scales. \\item  In two sources, SDSS J1333+0012 and SDSS J2215-0045, there is appearance/disappearance of BAL components.  The observed frequency of 9\\% for BAL appearance/disappearance is different from the value, 3\\% found for \\civ\\ BALs \\citep[][]{filiz12}. Confirming this possible difference is important for understanding the different BAL QSO classes  in the frame work of  \"orientation\" or \"evolutionary\" scenario. \\item As has been reported for \\civ\\ BALs,  we also find a trend for  larger fractional change in absorption equivalent widths to occur  in shallower and higher velocity features. We do not find any correlation between BAL variations and the QSO luminosity. However, it is interesting to note that there is a higher  fraction of objects in the high velocity high luminosity range. \\item The observed fraction of sources with varying BALs in  FeLoBAL QSOs (i.e, 20\\%) is small compared to  LoBAL QSOs (44\\%) and HiBALs (67\\%).  Our observation suggests that   FeLoBAL QSOs could be a different population compared to  LoBAL QSOs not showing \\feii\\ absorption. Such a difference is expected based on dust emission and distance estimates of clouds from photoionization models. However, as the predictions for the variability time-scales are not available, we are not in a position to distinguish between the two proposed scenarios of  FeLoBAL QSOs. \\item We find a larger fraction of sources with varying BALs in a sub-sample of Fe emitting sources. \\citet{zhang10} have found that the  LoBAL QSO fraction is higher when the strength of \\feii\\ emission is high using SDSS-DR5 data. \\citet{xiao11} have found a strong correlation between \\feii\\ emission strength and Eddington ratio. Confirming the higher variability detection in BAL QSOs with \\feii\\ dominated continuum in a larger sample will establish the connection between Eddington ratio (i.e. accretion efficiency) and line variability. \\item We do not find any trend in the variation of the \\mgii\\ and \\aliii\\ BAL equivalent widths with the continuum parameters. This is consistent with what is seen in  previous studies on \\civ\\ BAL systems.  Wide range in equivalent width differences, $\\Delta$W,  seen for a given magnitude difference, $\\Delta$m and vice versa imply that absorption line variations are not dominantly driven by photo ionization changes induced by continuum variations.  Emergence and profile variations of BALs in some sources imply that  dynamical changes like movement of clouds across the line of sight, may be important. \\end{enumerate}"
    },
    "1402/1402.7088_arXiv.txt": {
        "abstract": "We present the results of a semianalytical model that evolves the masses and spins of massive black holes together with the properties of their host galaxies along the cosmic history. As a consistency check, our model broadly reproduces a number of observations, e.g. the cosmic star formation history, the black hole mass and luminosity function and the galaxy mass function at low redshift, the black hole to bulge mass relation, and the morphological distribution at low redshift. For the first time in a semianalytical investigation, we relax the simplifying assumptions of perfect coherency or perfect isotropy of the gas fueling the black holes. The dynamics of gas is instead linked to the morphological properties of the host galaxies, resulting in different spin distributions for black holes hosted in different galaxy types.  We compare our results with the observed sample of spin measurements obtained through broad K$\\alpha$ iron line fitting. The observational data disfavor both accretion along a fixed direction and isotropic fueling. Conversely, when the properties of the accretion flow are anchored to the kinematics of the host galaxy, we obtain a good match between theoretical expectations and observations. A mixture of coherent accretion and phases of activity in which the gas dynamics is similar to that of the stars in bulges (i.e., with a significant velocity dispersion superimposed to a net rotation) best describes the data, adding further evidence in support to the coevolution of massive black holes and their hosts. ",
        "introduction": "Observational evidence for the existence of massive black holes (MBHs) in the center of galaxies is circumstantial but convincing. Stellar orbits near the Galactic center point at the presence of a dark and massive object, which is too heavy and compact to be a cluster of low-luminosity bodies~\\citep{darkcluster} or a fermion star~\\citep{schoedel}. Also, near-infrared observations of this object show that it is likely to have an event horizon, because if it did not, it would emit thermal radiation and be more luminous than observed~\\citep{horizon1,horizon2}. This disfavors exotic horizon-less alternatives to General Relativity's black holes, such as e.g. boson stars and gravastars. Similarly, MBHs are expected to be present in the center of most galaxies (and not just ours), as their presence is required to explain quasars (QSOs), active galactic nuclei (AGNs)~\\citep{soltan82} and the cosmic downsizing~\\citep{AGN_downsizing1,AGN_downsizing2,AGN_downsizing3}. Indeed, observations of nuclear stellar and gas dynamics, reverberation mapping, spectroscopic single epoch measurements and SED fitting \\citep[see e.g.][and references therein for a discussion of the different methods]{ghez08, gillessen09, kormendy95, miyoshi95, hicks08, blandford82, kaspi05, calderone13, castignani13} corroborate this expectation and allow the mass of MBHs to be estimated. Furthermore, these MBH mass measurements have in turn resulted into the discovery of the MBH-galaxy relations \\citep[e.g.][]{gebhardt00, ferrarese00, marconi03, haring04, gultekin09}, hinting at a symbiotic evolution of the central compact object and its host galaxy.  MBHs are also expected to possess a spin angular momentum. The combination of mass and spin completely characterizes these objects if they are described by the black-hole solutions of General Relativity~\\citep{kerr}\\footnote{Black holes in General Relativity can also have an electric charge~\\citep{kerr_newman}, but that is expected to be quickly canceled by the charges in the plasma surrounding astrophysical black holes, as well as by quantum effects such as Schwinger pair production~\\citep{Gibbons:1975kk,Hanni:1982}, or vacuum breakdown mechanisms triggering cascades of electron-positron pairs~\\citep{1969ApJ...157..869G,1975ApJ...196...51R,blandford1977}.}. If instead the gravity theory describing our Universe is not exactly given by General Relativity, other charges may be needed to describe MBHs, but the spin will still be one of them ~\\citep[c.f. ][for a few examples of black holes in alternative gravity theories]{AE_BHs,CS_BHs,Pani_BHs,GB_BHs}. Indeed, future spaced-based gravitational-wave interferometers such as ESA's L3 experiment eLISA~\\citep{lisa13} will measure MBH masses and spins with fantastic accuracy (respectively $\\sim 0.1$\\% and $\\sim 1$\\%), and also put General Relativity to the test by e.g. measuring possible additional black-hole charges. More importantly from an astrophysical point of view, measurements of MBH spins could unveil still unknown links between the MBHs and their galactic hosts. As an example, knowledge of the spins could constrain the properties of the gas accreting onto the MBHs, e.g. whether it spirals toward it on a preferential plane, on completely isotropic directions, or with other more complex dynamics \\citep[][hereinafter D13]{dotti13}.  Until eLISA or a similar space-based gravitational-wave interferometer is launched, however, measuring the spins of MBHs is considerably more difficult than measuring their masses. This is because the spin affects the dynamics of stars and gas surrounding the MBH only  at distances orders of magnitudes smaller than the MBH influence radius \\citep[e.g.][]{bardeen75}. Currently, the most accurate way of measuring the spins of MBHs is through the spectra of relativistically broadened K$\\alpha$ iron lines \\citep[see][for recent reviews on the topic]{reynolds13, brenneman13}. Indeed, using current X-ray spectrographs (in particular XMM and Suzaku), a significant sample of MBHs with spin estimates is being built up. Also, more stringent constraints on the spin estimates are becoming available from hard X-ray data from NuSTAR \\citep{risaliti13,marinucci14a,marinucci14b}. Finally, spin measurements with K$\\alpha$ iron lines are one of the main scientific drivers of ESA's upcoming L2 mission ATHENA+ \\citep{nandra13}.  This situation calls for a theoretical framework capable of interpreting available data for MBH spins and making testable predictions for future measurements. However, detailed modeling of the MBH spin evolution is still missing. Since the spin is a vector, any theoretical model should take the evolution of its direction into account, because that also affects the evolution of the magnitude. This is particularly clear when one considers MBH mergers, which are expected in the late stages of galaxy mergers \\citep[e.g.][]{begelman80}. The final spin, mass and kick velocity of the resulting MBH remnant  critically depend on the orientations of the spins of the progenitors \\citep{spin_formula1,spin_formula2,BKL,mass_formula,kick_RIT,kick_goddard}.  The main driver of the MBH evolution, however, is not expected to be given by mergers but rather by radiative efficient gas accretion \\citep{soltan82, merloni04, shankar04, shankar13b, wang, li}. In this case, the spin magnitude evolution depends on whether the gas in the very central regions of the accretion disk is co-rotating or counter-rotating relative to the MBH spin. This, in turn, depends on the larger-scale geometry of the fueling process (i.e. if the gas keeps falling toward the MBH on some preferential plane), as well as on the possible re-alignment of the MBH-disk system due to the Bardeen-Petterson effect \\citep{bardeen75}. The latter is the interplay between the frame dragging from the rotating spacetime geometry and the viscous stresses in the disk. The details of the effect of the spacetime's geometry on the properties of accretion disks are still not completely understood \\citep[see, e.g.][and references therein]{sorathia13a, sorathia13b}. All the studies that follow the coupled evolution of the magnitude and direction of the MBH spin have been performed under some simplifying hypotheses: i.e. assuming that the disk has isotropic viscosity \\citep[see the discussion in][]{sorathia13a, sorathia13b}, using order of magnitude estimates for the effect of the alignment \\citep[e.g.][]{king06}, or under the small-deformation approximation for the warped disk \\citep[e.g.][]{perego09}.  These simplifying hypotheses allowed for simple recipes for the spin evolution that have been used in semi-analytical models of structure formation and evolution \\cite[see, e.g.][]{volonteri05, berti08, lagos2009, fanidakis11, fanidakis2, mymodel,volonteri_sikora, volonteri13}. However, most of these studies~\\citep{volonteri05, berti08, lagos2009, fanidakis11, fanidakis2} assume that the MBH fueling is either always coherent (i.e. accretion always occurs exactly on the equatorial plane), or statistically isotropic (i.e. the MBH accretes small gas clouds with randomly oriented angular momenta). These assumptions were partly relaxed in \\cite{mymodel}, \\cite{volonteri_sikora} and \\cite{volonteri13}, who noted that the Bardeen-Petterson effect is likely to align the MBH spin with the angular momentum of the accreting gas in gas-rich nuclear environments, thus making the accretion flow effectively coherent in such situations, as opposed to isotropic accretion, which is more likely in gas-poor nuclear environments.  Even with this improvement, however, the possible parameter space of the accretion flow is still limited to two points (coherent accretion in gas-rich environments vs isotropic one in gas-poor environments).  Recently D13 demonstrated that the spin evolution can be significantly different if a more realistic description of the accretion flow is adopted, namely one that allows regimes intermediate between a perfectly coherent and a perfectly isotropic flow. \\cite{dubois13} studied the evolution of the spins of a population of MBHs taking into consideration the dynamical properties of the fueling gas. They extracted the direction of the angular momentum of the accreting matter on the smallest scales resolved by hydrodynamical simulations (down to $\\approx 10$ pc in their highest resolution run). In a similar spirit, in this paper we couple the semianalytic model of  \\citet[][hereinafter B12]{mymodel},  to the spin evolution model of D13. In particular, we link the degree of coherency/isotropy of the accretion flow to the dynamical parameters of the gas and stars observed in galactic nuclei. We stress that, also in our case, most of the observations that we use have a spatial resolution $\\gsim 100$ pc, significantly larger than typical accretion disks scales, and that the dynamical properties of the gas could change at smaller unresolved scales \\citep[e.g.][]{hopkins12, maio13, dubois14}. In this sense, therefore, a comparison of the predictions of our model for the MBH spin evolution with actual spin measurements provides a way of testing the dynamics and coherency of the nuclear gas at yet unresolved scales. We attempt such a test with the MBH spin measurements from K$\\alpha$ iron lines available to date. We note, however, that our model provides a flexible framework to also interpret future, more accurate measurements from experiments such as eLISA and ATHENA+.  The paper is organized as follows. In Section \\ref{sec:model} we present the details of our model for the evolution of MBHs in their galactic hosts. Improving on B12, we adopt a different star-formation law (see Section \\ref{sec:SF} and Appendix \\ref{ap:SF}) and, more importantly, we update the prescription for the evolution of the MBH spins under accretion, following D13 (see Section \\ref{sec:spin}). The degree of coherence of the  accretion flow is then linked to the kinematic properties of the host galaxy in Section \\ref{sec:link}. We consider in particular three models: one where the degree of coherence is related to the gas kinematics, one that links it to the stellar kinematics, and a hybrid model where information from both the gas and stellar kinematics is used. In Section \\ref{sec:cali}, we calibrate the free parameters of our semianalytical galaxy-formation model against observations. We present our predictions for the evolution of MBH spins in Section \\ref{sec:results}, and in Section \\ref{sec:obs} we compare them to measurements from K$\\alpha$ iron line profiles. In Section \\ref{sec:conclusions} we draw our conclusions.  \\begin{figure} \\includegraphics[width=8.5cm]{model.pdf} \\caption{\\label{fig:model} Schematic representation of our semianalytical model, based on that of B12. Note that B12 did not use the distinction between ``bulges'' forming from major galactic mergers and ``pseudobulges'' forming from bar-instabilities of disk galaxies (although both formation channels were present). In this paper, instead, we assume that bulges and pseudobulges have different star-formation laws (c.f. Section \\ref{sec:SF}) and, in some realizations of our model, different accretion properties onto the central MBH (c.f. Section \\ref{sec:link}). The other major change from the model of B12 is indeed the prescription for the spin evolution of MBHs under accretion and mergers, which is schematically represent in Figure~\\ref{fig:flow} and described in detail in Section \\ref{sec:spin}.} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we presented the results of a semianalytical model for the evolution of galaxies and MBHs. The model keeps track (although in simplified ways) of the morphology of the galaxies, as well as the MBH masses and spins. For the first time we link the dynamical properties of the gas fueling the MBHs to the host galaxy properties, through different observationally based prescriptions, and derive predictions testable with existing observations. We stress that all the other similar investigations predicting MBH spin distributions either assumed ``chaotic'' accretion (MBHs accreting small gas clouds with isotropically oriented angular momenta) or ``coherent'' accretion (MBHs accreting all the time on a fixed plane).  Our model predicts different spin distributions for different types of galaxies. To date, only $\\approx 20$ MBH spins have been directly measured through K$\\alpha$ iron line fitting \\citep{reynolds13, brenneman13}. All these MBHs (but one) are hosted in low-redshift late type galaxies, preventing us from testing our results for different galaxy types. We therefore tested our model by selecting low-redshift accreting MBHs hosted in spirals. { For this galaxy class, we have assumed three different gas dynamics prescriptions. In the {\\it pseudobulge} model we assumed that the dynamics of gas fueling the central MBHs is similar to that of the stars both in galaxy  bulges or pseudobulges, according to the nature of the host's spheroidal component. In the {\\it disk} model, instead, the fueling gas has the more coherent dynamics of the large scale gaseous disk. The {\\it hybrid} model shares the same gas dynamics as the {\\it disk} model for isolated galaxies, but assumes a gas dynamics similar to that of stars in spheroids during merger driven accretion events (see Section~\\ref{sec:link} for details)}.  Our models predict interesting features in the MBH spin distribution, and a statistical comparison to observations yields a number of interesting results: \\begin{enumerate} \\item different galaxy morphologies result in different MBH spin distributions. Because the geometrical properties of the accretion flow are related to the large scale kinematics of the galaxy, MBHs hosted in ellipticals tend to have lower spins than those hosted in spirals; \\item in general, MBH spins are a decreasing function of the MBH mass: $\\sim 10^{6}\\msun$ MBHs tend to be maximally spinning, whereas for masses $>10^{8}\\msun$ a wide range of values is possible for the spin, depending on the host morphology and MBH accretion rate; \\item accreting MBHs in spirals, i.e. those matching the observational sample, tend to spin fast, and do {\\it not} provide an unbiased indicator of the underlying spin distribution of the overall MBH population; \\item 2D-KS tests show that, in general, the {\\it hybrid} and {\\it pseudobulge} models are consistent with observations, while the {\\it disk} model is disfavored. A quantitative assessment of the compatibility of the {\\it disk} model with the data, however, strongly depends on the maximum value to which MBHs can be spun up by accretion. This issue is discussed in detail in Appendix \\ref{app:sanity}, where we show that if this value is $a_{\\rm bh}=0.998$~\\citep{thorne}, the {\\it disk} model can be safely ruled out at $>99\\%$ confidence level; \\item a likelihood odds ratio study shows that the {\\it hybrid} model -- which takes into account the possible angular momentum reshuffling due to  merger-driven phases of nuclear activity --  provides the best match to the observed data. In a two model comparison: $i)$ it is strongly favored (at $>99\\%$ confidence level) over the {\\it disk} model; $ii)$ it is moderately favored (only at about $90\\%$ confidence level) over the {\\it pseudobulge} model; \\item For the sake of comparison, in Appendix \\ref{app:co-cha} we have also run three separate sets of simulations adopting the standard coherent-accretion model, as well as two flavors of the chaotic-accretion model. These models provide a much worse description of the data; the {\\it coherent} model shows features similar to the {\\it disk} model, but with a poorer agreement to the data, while the chaotic models are unambiguously ruled out by observations. \\end{enumerate}  We emphasize that the study of the evolution of the MBH spin distribution in a cosmological framework is still in its infancy. Because our model is idealized from many points of view, here we highlight some still-needed theoretical and observational improvements. $i)$ The treatment of the Bardeen-Petterson effect, which plays a major role in determining the fraction of prograde vs retrograde accretion events, is based on simple isotropic-$\\alpha$-viscosity driven accretion disks, and this assumption has been criticized by \\cite{sorathia13a, sorathia13b}. The impact of this strong and debated assumption is unknown, since no working models of the MBH spin evolution relaxing it are available to date. $ii)$ We assumed that every gas cloud fueling the MBHs weighs $M_{\\rm cloud}\\sim 3\\times10^4 \\msun$. Giant molecular clouds are observed to have a broad mass spectrum. Observational studies of local galaxies reported typical molecular cloud mass functions with cut-offs at about $10^{5 - 6} \\msun$ \\citep[e.g.][]{williams97, engargiola03, fukui08}, making the average cloud mass comparable to the mass we that assume in this paper. We note however that the cut-off mass seems to depend on the host galaxy properties \\citep[e.g.][and references therein]{hopkins12b}. A simple implementation of a ``characteristic'' cloud mass function is not trivial, and beyond the scope of this paper. The effect of different average cloud masses can anyway be easily understood: the cloud mass contributes to setting the maximum MBH mass at which a complete alignment of the spin occurs at every accretion event, effectively making the fueling dynamics unimportant. To first approximation (see D13 for a more detailed discussion), an order of magnitude variation of the average cloud mass results in an order of magnitude shift along the MBH mass axis of all the spin distributions. The results not showing any significant feature in the spin evolution (e.g. the $f_{\\rm Edd}>0.01$ {\\it disk} model at low redshift, but see also the {\\it coherent} and {\\it chaotic I} models discussed in Appendix B) would thus remain mostly unchanged under a variation of the cloud mass. On the other hand, the {\\it pseudobulge} and {\\it hybrid} models would predict higher spins for higher masses (lower spins for lower masses) if the cloud mass were higher (lower). We checked that a shift of the  spin distributions by one order of magnitude toward larger masses does not affect the compatibility of those two models with the data, yielding a KS probability $p_{\\rm KS}>0.3$. A similar shift of the distribution toward lower masses results in a slightly worse description of the data, but still with a KS probability $\\gsim 0.1$. $iii)$ Our dynamical prescriptions for the nuclear gas dynamics are mostly based on gas and/or stellar dynamics observations on scales much larger than the accretion disk or the MBH sphere of influence. Our model would greatly benefit from a description of the gas dynamics on scales $\\lsim 100$ pc, already achievable in some cases with high resolution ALMA spectroscopy \\citep[e.g.][]{combes13, combes14}. Finally, $iv)$ in our model we completely neglected the effect of rotational energy extraction from the MBH on the spin. This could in principle power relativistic jets through the Blandford-Znajek mechanism, limiting the maximum spin to lower values in radio galaxies.  From an observational point of view, the main improvement would consists in enlarging the currently small sample of measured MBH spins, and to extend it to higher redshift and different galaxy types{\\footnote{{ While finalizing this manuscript, \\cite{reis14} reported a spin measurement of the $z=0.658$ lensed quasar 1RXS J113151.6-12315 of $a=0.87^{+0.08}_{-0.15}$ (3$\\sigma$ errors). The estimated MBH mass is $2\\times10^{8}\\msun$, the accretion rate computed from its bolometric luminosity is nearly Eddington, and the host is a Seyfert 1 spiral galaxy \\citep{claesken06}. Thus, this system falls in our sample of accreting spirals at $z=1$ shown in figure \\ref{spinobs}, and its inferred mass and spins are perfectly consistent with all the models shown there.}}}. In this regard, the forthcoming Astro-H satellite will have exquisite combined spectral resolution and sensitivity, making K$\\alpha$-based spin measurements possible perhaps up to $z\\sim1$ \\citep{astroH2012}. Along the same lines, Athena+, to be launched in 2028, will push such measurements to $z\\sim2$ for extremely X-ray bright systems, significantly expanding the observed spin sample \\citep{dovciak2013}. On a shorter timescale, the eROSITA satellite will observe $>10^5$ AGNs. Image stacking in luminosity and redshift bins will reveal the 'average' shape of the K$\\alpha$ line of these systems, making possible to study trends in the typical spin of AGNs across a large mass range  and  for redshifts perhaps up to $z\\sim1$~\\citep{merloni2012}, thus providing an important benchmark for comparison to theoretical models. On a different note, the eLISA mission, now selected by ESA for the L3 launch slot, will measure the spins of merging MBH binaries across the cosmic history to high precision, providing a sample that will not be biased toward high X-ray luminosities.  Even though the future of spin measurements looks literally bright, we should also keep in mind that the uncertainties on the value of measured spins depend on the technique used. Different methods have been used to constrain the MBH spin distribution at different redshifts and for different masses. Some methods estimate the spin value from the radio properties of the AGNs \\citep[e.g.][and references therein]{dali13}, from the ionizing flux required to produce the observed broad emission lines \\citep{netzer13}, or fitting the continuum from the accretion disk \\citep{laordavis11}. Some of these methods are still matter of debate \\citep[e.g.][]{gallo12, raimundo12, laor11}, and all of them require an independent measure of the MBH mass to estimate its spin. In this paper we compared our predictions with the mass independent measurements obtained through relativistically broadened iron K$\\alpha$ line fitting, to prevent any possible spurious mass spin correlation. We refer to \\cite{reynolds13} and \\cite{brenneman13} for a detailed discussion of the uncertainties of these measurements. Since our models predict a significant dependence of the MBH spins on the MBH masses, the uncertainties on the masses are fundamental as well. For { example, Mrk 359} has an estimated mass of $\\approx 10^6 \\msun$, obtained through a single epoch measurement of the continuum and broad H$\\beta$ FWHM \\citep{2005ApJ...618L..83Z}. The width of the line is only 480 km s$^{-1}$ significantly smaller than the typical value for type 1 AGN as well as than the threshold for being classified as a NLS1. Since the MBH is hosted in a spiral galaxy the line is peculiarly large to be associated with a single narrow line, and a broad component is expected because of the point like emission clearly visible in the HST imaging \\citep{1998ApJS..117...25M}. The observed line width could be so narrow because of orientation effects, already suggested to be significant for NLS1s \\citep[][]{decarli08}, causing an underestimation of the MBH mass. (To be on the safe side, we checked that the contribution of Mrk 359 to our analysis is negligible. Its MBH spin is so poorly constrained that removing it from the observational sample has no effect on our results.) More independent measurements of the masses of MBHs with spin estimates would help better constrain their values. The detection of gravitational waves will also greatly improve the situation, giving high precision measurements of both the MBH masses and spins, which could be compared with predictions for merging systems.  { Regardless of the large uncertainties on which we commented above, it is remarkable that two of the models ({\\it pseudobulge} and {\\it hybrid}) that we investigated generate samples of MBH masses and spins consistent with observations. Since these models assume that the accreting gas has a less coherent dynamics than the larger-scale gas structure (although, in the hybrid case, only in recent merger remnants)}, our analysis seems to suggest the existence of a physical process that, while decreasing the gas angular momentum magnitude (triggering accretion in the first place), reshuffles the gas angular momentum direction as well.  Such a reshuffling could be caused by local processes, i.e. star formation and the torques exerted by the gas self-gravity \\citep[e.g.][]{maio13}, as well as larger-scale gravitational instabilities and violent gas inflows \\citep[e.g.][]{hopkins12, dubois14}.  On the other hand, in order to achieve such a good agreement between the predictions of the model and the observed spins, we have to require the accreting gas to have, on average, a non-zero angular momentum, i.e. the gas must \\textit{not} accrete isotropically onto the MBH. We have in fact shown that a perfectly isotropic accretion flow would result in significantly lower spins for intermediate to large MBH masses, which is inconsistent with observations (c.f. Appendix \\ref{app:co-cha}).   { We further note that the {\\it pseudobulge} and the {\\it hybrid} models, while providing the best descriptions of the data, hint at different (although possibly coexisting) evolutionary scenarios. The hybrid scenario has the highest statistical significance and a simple physical interpretation, with a clear candidate for the source of turbulence.  The gas dynamics remains quite coherent for most of the time, which is reasonable because gas can efficiently dissipate turbulent motions on few local orbital times \\citep[see the discussion in][and references therein]{maio13}. Only during galaxy mergers the violent reshuffling of the angular momentum allows for more isotropic gas inflows. We note that, strictly speaking, in the {\\it hybrid} model we adopt a ``bulge-like'' dynamics for the accreting gas during the {\\it whole} of the merger driven accretion events, even though the AGN activity lasts for more than the galaxy merger itself. Thus, during the tails of the nuclear activity, accretion may be more coherent than in our model's assumptions. However, during a merger-driven AGN phase most of the gas is accreted in a short almost-Eddington burst, lasting about a Salpeter time $\\lesssim 0.2$ Gyr.  The long lived accretion tails, still observable at low redshifts,   present instead a lower Eddington ratio $f_{\\rm Edd}\\sim 0.01$ and cannot therefore significantly change the MBH spin magnitude.  Although the pseudobulge model is disfavored by our statistical analysis, it still provides a good description of the observational data and is consistent with the evidence that low redshift AGNs with moderate to high Eddington ratios are preferentially hosted in pseudobulges \\citep{heckman14}. Pseudobulges have recently been associated with stellar bars, i.e., structures able to drive gas inflows toward the galaxy nucleus \\citep[see, e.g.][for up to date reviews]{sellwood13, kormendy13}. After the initial instability responsible for the bar formation, the bar itself generally undergoes a thickening that can be either caused by a second ``buckling'' instability or by a vertical resonance~\\citep{sellwood13}, possibly responsible for the pseudobulge formation~\\citep{kormendy13}. In this scenario, the gas could accrete with a more disk like dynamics until the bar thickens, and fall toward the MBH with a higher velocity dispersion after the pseudobulge forms if the buckling instability and/or any vertical resonances significantly affect the gas motion.  The paucity of data and the many uncertainties about the different fueling processes do not allow us to firmly reject either of these models.  If either model is confirmed by more data, it could point toward a physically motivated galaxy-MBH coevolution scenario for low redshift disk galaxies, in which the dynamics of gas in nuclear bars and/or mergers drive the MBH mass and spin evolution.}"
    },
    "1402/1402.0661_arXiv.txt": {
        "abstract": "The AllWISE processing pipeline has measured motions for all objects detected on WISE images taken between 2010 January and 2011 February. In this paper, we discuss new capabilities made to the software pipeline in order to make motion measurements possible, and we characterize the resulting data products for use by future researchers. Using a stringent set of selection criteria, we find 22,445 objects that have significant AllWISE motions, of which 3,525 have motions that can be independently confirmed from earlier 2MASS images yet lack any published motions in SIMBAD. Another 58 sources lack 2MASS counterparts and are presented as motion candidates only. Limited spectroscopic follow-up of this list has already revealed eight new L subdwarfs. These may provide the first hints of a ``subdwarf gap'' at mid-L types that would indicate the break between the stellar and substellar populations at low metallicities (i.e., old ages). Another object in the motion list -- WISEA J154045.67$-$510139.3 -- is a bright ($J \\approx 9$ mag) object of type M6; both the spectrophotometric distance and a crude preliminary parallax place it $\\sim$6 pc from the Sun. We also compare our list of motion objects to the recently published list of 762 WISE motion objects from \\cite{luhman2014}. While these first large motion studies with WISE data have been very successful in revealing previously overlooked nearby dwarfs, both studies missed objects that the other found, demonstrating that many other nearby objects likely await discovery in the AllWISE data products. ",
        "introduction": "The Wide-field Infrared Survey Explorer (WISE; \\citealt{wright2010}) was built to survey the sky simultaneously in four bands with central wavelengths of 3.4, 4.6, 12, and 22 $\\mu$m (hereafter referred to as W1, W2, W3, and W4, respectively). The four-band cryogenic mission covered the sky 1.2 times between the dates 2010 Jan 07 and 2010 Aug 06. After the solid hydrogen cryogen was depleted in the outer tank, data covering 30\\% of the sky continued to be acquired in the three-band (W1, W2, and W3) cryogenic mission until 2010 Sep 29, when cryogen in the inner tank was exhausted. WISE surveyed 70\\% of the sky in the two-band (W1 and W2) post-cryogenic NEOWISE mission (\\citealt{mainzer2011}) until the satellite was placed into hibernation on 2011 Feb 01. In total, WISE surveyed the full sky twice in at least two bands and 20\\% of the sky a third time, with each coverage epoch separated by approximately six months.  Data were processed and released separately for the cryogenic and post-cryogenic missions. The aim of the AllWISE program was to combine all of the data from the WISE mission to leverage the full depth and time history of the W1 and W2 data. Specifically, one of these goals was to use the full WISE data to perform the first all-sky proper motion survey at these wavelengths. In addition to the exciting prospects created by doing a kinematic survey in a new wavelength regime, AllWISE processing -- because of the short, six-month time baseline between epochs -- enables us to search for very high motion objects that may have been overlooked by the longer time baselines of earlier all-sky motion surveys. In these surveys, an object could be ``lost'' between consecutive views. A case in point is the first motion discovery from AllWISE, the usdM3 star WISEA J070720.50+170532.7 (\\citealt{wright2013}), which is moving at $1\\farcs8$ yr$^{-1}$ and was missed by all previous motion surveys despite being in the more heavily studied northern hemisphere, being easily detected at optical wavelengths, and falling in an uncomplicated region free of source confusion.  In this paper, we describe how motions are measured in AllWISE (section 2), characterize the resulting motions and present caveats to users of the AllWISE data products (section 3), present a catalog of new motion discoveries (section 4), compare that to the recent list of WISE motion discoveries published by \\cite{luhman2014} (section 5), discuss our spectroscopic follow-up of some of the discoveries (section 6), and highlight a couple of specific science cases -- the hunt for rare, L-type subdwarfs (section 7) and the search for missing members of the immediate Solar Neighborhood (section 8) -- that we have already explored using the discovery list. In section 9 we provide a summary of our AllWISE results thus far. ",
        "conclusions": "We have characterized the motion measurements contained within the AllWISE Data Release. We have presented a list of 3,525 motion objects lacking previous literature in SIMBAD and find two very interesting results. The first is that AllWISE is revealing many more examples of low-metallicity objects near the stellar/substellar break, allowing researchers to probe the region near the hydrogen-burning minimum mass for objects formed early in the history of the Milky Way. The second is that AllWISE is continuing the WISE legacy of uncovering previously overlooked, very nearby objects. Prior to AllWISE processing, WISE had already revealed several nearby systems never before recognized --e.g., the L+T binary WISE J104915.57-531906.1 at 2.0 pc (\\citealt{luhman2013}) and the T6 dwarf WISEPC J150649.97+702736.0 (\\citealt{kirkpatrick2011}) possibly as close as 3.4 pc (\\citealt{marsh2013}); AllWISE has already contributed another overlooked object -- WISEA J154045.67$-$510139.3 ($\\sim$5.9 pc) -- to this distinguished list of WISE discoveries within the canonical 8-pc sample. The motion measurements provided by AllWISE will likely provide researchers many new gifts over the years to come."
    },
    "1402/1402.5791_arXiv.txt": {
        "abstract": "The steady state propagation of a phase transition front is classified, according to hydrodynamics, as a deflagration or a detonation, depending on its velocity with respect to the fluid. These propagation modes are further divided into three types, namely, weak, Jouguet, and strong solutions, according to their disturbance of the fluid. However, some of these hydrodynamic modes will not be realized in a phase transition. One particular cause is the presence of instabilities. In this work we study the linear stability of weak detonations, which are generally believed to be stable. After discussing in detail the weak detonation solution, we consider small perturbations of the interface and the fluid configuration. When the balance between the driving and friction forces is taken into account, it turns out that there are actually two different kinds of weak detonations, which behave very differently as functions of the parameters. We show that  the branch of stronger weak detonations are unstable, except very close to the Jouguet point, where our approach breaks down. ",
        "introduction": "In a cosmological first-order phase transition, bubbles of the stable phase nucleate and expand in the supercooled metastable phase. The motion and collisions of bubble walls are associated with the formation of cosmological remnants, such as gravitational waves \\cite{gw}, magnetic fields \\cite{gr01}, topological defects \\cite{vs94}, a baryon asymmetry \\cite{ckn93}, or baryon inhomogeneities \\cite{w84,h95,ma05}. The propagation of bubble walls is driven essentially by the pressure difference between the two phases, but is significantly affected by hydrodynamics \\cite{hidro} (for recent studies, see \\cite{ms09,ekns10,lm11,kn11,ms12}). Indeed, besides stirring the fluid, these phase transition fronts reheat the plasma, due to the release of latent heat. The back-reaction of such disturbances hinders the wall motion. In addition, the microscopic interaction of particles of the plasma with the wall causes, at the macroscopic level, a friction force \\cite{micro,bm09,hs13,ariel13}. In general, after a time which is much shorter than the total duration of the phase transition, the wall reaches a steady state with constant velocity (however, the wall may also run away \\cite{bm09}).  As a consequence of nonlinear hydrodynamics, different kinds of stationary solutions exist. Thus, the phase transition front may propagate, in principle, either as a weak, Jouguet, or strong deflagration, as well as a weak, Jouguet, or strong detonation. The latter, however, is not possible, since its fluid profile cannot fulfil the boundary conditions. These various propagation modes coexist in some ranges of parameters, and it is not easy to determine, in general, which of them will be actually reached in the phase transition. It will certainly depend on the initial and boundary conditions during the early transitory stage.  In the literature, the wall velocity is often given convenient values or left as a free parameter. For instance, slow-moving weak deflagrations are often assumed for baryogenesis, while fast-moving detonations are assumed for gravity-wave generation. In the latter case, a Jouguet detonation is often assumed. This solution corresponds to the lower bound for the detonation velocity. The choice of this particular solution avoids calculating the wall velocity, which would involve considering the balance of the driving and friction forces. This assumption is motivated by the case of chemical burning, where weak detonations are forbidden  \\cite{s82}. However, a phase transition is qualitatively different from chemical burning, and weak detonations are possible in addition to Jouguet detonations \\cite{l94}.  Even when a complete calculation of the wall velocity is performed, there still remains the problem of the existence of multiple solutions (for a recent discussion, see \\cite{ms12}). Besides the coexistence of the aforementioned kinds of solutions,  a calculation of the wall velocity may give double-valued solutions of a given kind. In particular, for the weak detonation case, there are in general two branches. One of these branches consists of higher velocity solutions, which are hydrodynamically weaker (they cause smaller disturbances of the fluid). These solutions approach the speed of light for low enough friction or high enough supercooling. The other branch of weak detonations consists of stronger solutions with lower velocities, and ends at the Jouguet point. These detonations behave unphysically as a function of the parameters. Therefore, weak detonations might be divided into ``weaker'' weak detonations and ``stronger'' weak detonations.  An important issue is, thus, to determine the actual propagation mode for a given set of parameters. A stability analysis provides a useful tool to determine the final state of the wall motion, since not all of the possible hydrodynamic solutions turn out to be stable. For instance, it is known that strong deflagrations are generally unstable, and that weak deflagrations are unstable in certain ranges of parameters \\cite{hkllm,stabdefla}. The case of detonations is less clear.  In the cosmological context, the stability of weak detonations was first discussed by Huet et al. \\cite{hkllm}. The standard approach is to consider small perturbations of the fluid variables on each side of the wall, together with small deformations of the latter, which is assumed to be infinitely thin \\cite{landau,link}. The fluid perturbations on the two phases are linked by junction conditions at the interface. In the case of detonations, the fluid enters the wall supersonically (in the reference frame of the latter). As a consequence, fluid perturbations cannot evolve in front of the wall. This fact lead Huet et al. to conclude that perturbations cannot grow at all, and weak detonations are always stable. However, numerical calculations of the time evolution of the wall-fluid system seem to indicate that configurations belonging to the ``stronger'' branch of weak detonations are unstable \\cite{ikkl94}.  In fact, as noted by Abney \\cite{abney}, the conclusion of Ref. \\cite{hkllm} is incorrect, as fluid perturbations may vanish in front of the wall and not behind it (and fulfil the junction conditions). Such perturbations, which were not considered in \\cite{hkllm}, may be unstable. Unfortunately, Abney (and later Rezzola \\cite{r96}) considered only Jouguet and strong detonations. As already mentioned, the latter are not possible in a phase transition. Moreover, even in the case of Jouguet detonations, the treatment of Refs. \\cite{abney,r96} is not valid. In the first place, an equation for the interface (depending on the driving and friction forces) was not considered in these works, even though its importance had been already pointed out be Huet et al. \\cite{hkllm}. Most important, the treatment considers perturbations around a solution of constant fluid velocity $v$ and temperature $T$, whereas the fluid profile for a Jouguet detonation is that of a rarefaction wave behind the wall, with $v$ and $T$ varying in space and time.  In this paper we shall consider the hydrodynamic stability of detonation fronts in a cosmological phase transition. We shall consider only weak detonations. Unfortunately, for the case of Jouguet detonations, the fact that the fluid profiles are not constant behind the wall makes the treatment too involved. For the same reason, our approach will break down for weak detonations which are very close to the Jouguet point. Except in this limit, we will show that the lower velocity branch of weak detonations is unstable under linear perturbations at all wavelengths. Our approach is essentially the same we used in Ref. \\cite{stabdefla} for the case of weak deflagrations, and we shall refer to that work for some details. Due to the vanishing of perturbations in the metastable phase, the treatment of detonations turns out to be much simpler than that of deflagrations. This will allow us to obtain analytical and model-independent results.   The plan of the paper is the following. Before performing the stability analysis, we devote the next section to a detailed discussion on the structure of weak detonations. We also consider the coexistence with runaway solutions. In Sec. \\ref{staban} we solve the equations for the linear perturbations of the wall-fluid system, and in Sec. \\ref{stab} we discuss in detail the velocity intervals corresponding to unstable solutions and to the validity of our approach. We also discuss some implications of our results for a cosmological phase transition. Our conclusions are summarized in Sec. \\ref{conclu}. ",
        "conclusions": "\\label{conclu}  The possible cosmological consequences of a first-order phase transition depend on the dynamics associated to the motion of bubble walls. It is well known that several hydrodynamic solutions are in principle possible for the propagation of these phase-transition fronts. In particular, for the steady state motion we may have weak or Jouguet detonations, as well as weak, Jouguet or strong deflagrations. Nevertheless, some of these solutions are in fact unstable. The stability analysis is thus important in order to determine which of the propagation modes will be actually realized in a phase transition.  In this work we have studied the stability of detonations under small perturbations of the fluid and the interface. This paper is a sequel of a similar study for the case of deflagrations \\cite{stabdefla}. The treatment of detonations is quite simpler, due to the fact that the incoming fluid (in the reference frame of the wall) is supersonic. Thus, the fluid in front of the wall is not affected by the phase-transition front. This has two practical consequences for the calculation. On the one hand, for the stationary case, the incoming fluid velocity is given by the wall velocity, $v_+=-v_w$, and the temperature in front of the wall is given by the boundary condition $T_+=T_N$. On the other hand, perturbations originated at the wall cannot grow in front of it. This simplifies the treatment of the stability, since only perturbations of the outgoing fluid must be considered (i.e., perturbations which vanish in front of the wall) \\cite{abney}. These simplifications allowed us to obtain analytical and model-independent results.   Before the stability analysis, we have discussed in detail the detonation solutions. It is well known that strong detonations are hydrodynamically forbidden \\cite{s82}. Weak detonations can exist in the velocity range $v_J<v_w<1$, while the Jouguet detonation corresponds to the case $v_w=v_J$. The Jouguet velocity $v_J$ depends on the parameters of the equation of state, and is only determined by the temperature $T_N$. However, the precise value of $v_w$ within this interval will be determined by the balance of the driving and friction forces, and will further depend on friction parameters. Due to strong hydrodynamic effects, it turns out that some velocities in this range may not be realized. Besides, the resulting velocity turns out to be a double valued function of the parameters, and the bubble wall will need to ``choose'' which of the two possible hydrodynamic configurations to adopt.  Thus we have two branches of weak detonation solutions. One of them, with velocities in the range $v_{\\mathrm{crit}}<v_w<1$, where $v_{\\mathrm{crit}}$ is given by Eq. (\\ref{etamax}), corresponds to weaker weak detonations. The other branch has velocities in the range $v_J<v_w<v_{\\mathrm{crit}}$ and corresponds to stronger weak detonations. This branch behaves rather unphysically with the parameters. Although weak detonations are generally believed to be stable \\cite{hkllm}, numerical simulations \\cite{ikkl94} suggest that these stronger solutions are not.  We have applied the standard stability analysis for fluid interfaces \\cite{landau} to the case of relativistic detonation fronts in a phase transition. We have pointed out that this analysis breaks down at the Jouguet point. This is because the approach considers perturbations around constant velocity and temperature, which is not the case of the Jouguet detonation profile. As a consequence, previous results on the stability of Jouguet detonations \\cite{abney,r96} are not valid. We have shown that  weak detonations of the lower velocity branch are generally unstable under linear perturbations at all wavelengths. More specifically, weak detonations are unstable in the range $v_{\\mathrm{crit}}'<v_w<v_{\\mathrm{crit}}$, where $v_{\\mathrm{crit}}'$ is roughly given by Eq. (\\ref{vcritpriapp0}). Below $v_{\\mathrm{crit}}'$, our approach breaks down.  Nevertheless, as we have seen, $v_{\\mathrm{crit}}'$ is very close to the Jouguet velocity $v_J$. Therefore, it is quite unlikely that the actual value of the wall velocity (taking into account the friction) will fall in this very small interval. Regarding the higher velocity branch, the linear perturbation analysis is not conclusive, but we have argued that these solutions will probably be stable. This means that, in the case of coexistence of a detonation and a runaway solution, it is the stationary solution the one which will be realized.  Our main result is, thus, that the branch of weak detonations which are closer to the Jouguet point will not be realized during a phase transition. Unfortunately, these solutions are cosmologically interesting due to their strong disturbance of the fluid. We have discussed, in particular, how our results affect a mechanism of electroweak baryogenesis with detonation walls \\cite{cn12}. This mechanism is based on the fact that the reheating behind the wall may exceed the critical temperature, allowing the nucleation of symmetric-phase bubbles in the superheated fluid inside the broken-symmetry bubbles. The regions in parameter space for this scenario are significantly reduced once unstable detonations are discarded."
    },
    "1402/1402.1198.txt": {
        "abstract": "We searched the light curves of over 40000 stars in the 2 Micron All Sky Survey (2MASS) calibration database, spanning approximately 4 years, for objects that have significant day long dimming events. We also searched the multi-color light curves for red dimming events that could be due to transient extinction. In the color independent sigma-limited search, we found 46 previously unknown eclipsing binaries, 6 periodic variable stars %xxx likely to be intrinsic pulsators and 21 young stellar objects in the $\\rho$ Ophiucus star formation region previously studied by Parks et al. 2013. An additional 11 objects exhibited dimming events and most of these are unclassified. %Those without significant variation in near-infrared color  could be eclipsing binaries, %and two are active galaxies that exhibit correlations between color and brightness. % and we have failed to identify a period because we have not observed enough eclipses.  The search for red dimming events primarily reveals a population of low luminosity active galaxies that become bluer when they are brighter, and variable young stellar objects exhibiting high cross-correlation coefficients between color and brightness. The young stellar objects exhibit brightness and color variations in the direction of interstellar extinction whereas the active galaxies can have a bowed distribution in color and magnitude with reduced variation in color when the object is brightest. By plotting cross-correlation coefficients (between color and brightness) against the slope fit to points in the color/magnitude distribution, we find that active galaxies  tend to have lower slopes and correlation coefficients.  Among the objects that are usually quiescent (not strongly variable), we failed to find any dimming events deeper than 0.2 magnitude and lasting longer than a day. Two of the least embedded young stellar objects, however, dimmed by 0.2 mag for longer than a day without strong color variation. Having failed to find new exotic objects, we explored ways to eliminate commonly found objects so that a larger number of objects may be searched. We find that  all but one eclipsing binary are excluded by requiring moderate color variation during a dimming event and most of the active galaxies and young stellar objects are excluded by placing a limit on the standard deviation of the magnitude distribution in the light curve. ",
        "introduction": "One star in a young stellar binary system can host a circumstellar disc and the disc can periodically occult the other star as seen from a distance (e.g., $\\epsilon$ Aurigae, \\citealt{guinan02,kloppenborg10,chadima11}; EE-Cep, \\citealt{mikolajewski99,graczyk03,mikolajewski05,galan12}; OGLE-LMC-ECL-17782, \\citealt{graczyk11};  OGLE-LMC-ECL-11893, \\citealt{dong14}). Because discs can be large, the probability that a randomly oriented but young system exhibits eclipses may be fairly high \\citep{mamajek12,meng14}. Circumbinary disks can also periodically occult a star (e.g., KH 15D, \\citealt{kearns98,winn04}; V718 Per, \\citealt{grinin08}). Periodic and aperiodic occultations have been detected in young stellar objects \\citep{bouvier07,grinin08,moralescalderon11,parks13,wolk13,cody14}. Transient single, long, and deep occultation events of pre-main sequence stars have also been reported \\citep{rodriguez13,mamajek12}. Occulting disks are valuable as they would allow the structure of a distant disk to be studied on smaller spatial scales than possible with direct imaging and it would be possible to study their composition through absorption of star light.  The Two Micron All-Sky Survey (\\citealt{skrutskie06}; 2MASS) imaged the entire sky in three near infrared wavelength bands, $J, H, K$, between 1997 and 2001. Photometric calibration for 2MASS relied upon repeated observations of 35 calibration fields that are distributed at approximately 2 hour intervals in right ascension near declinations of 0 and $\\pm 30$ deg. %Each 2MASS calibration or tile covers a region 8.5\u00d5 wide (in right ascension) by 60\u00d5 long (in declination). % All source extractions from the calibration scans are available int the 2MASS Calibration Point Source Working Database (Cal-PSWDB). The database contains over 191 million source extractions. %derived from 73,230 scans Further descriptions of the Cal-PSWDB are available by \\citet{cutri06}, and \\citet{plavchan08} who found 247 red variable objects, including extragalactic variables and 23 periodic variable stars. Using one of the calibration fields, \\citet{parks13} recently presented a variability study of the $\\rho$ Ophiucus molecular cloud and star formation region, discovering 22 new cluster members via their near-infrared variability and periodicity in 35 young stellar objects. %Among the periodic variables %were a few M-type eclipsing binaries.  Calibration tiles were observed once per hour.  While numerous photometric monitoring projects search for objects that brighten, such as novae or micro lenses,  there have been few systematic searches for rare, individual dimming events. Single, isolated, dimming events are often ignored because they can be caused by poor photometric measurements that occur because of seeing variations, scintillation, clouds, imaging artifacts, and source confusion. Variability studies  search for periodic phenomena, including planetary transits, (e.g., \\citealt{colliercameron06,palaversa13}), use the moments of the magnitude distribution to find and classify variables (e.g., \\citealt{ciocca13,graczyk10}) or combine periods found in one band with color or spectral information (not necessarily as a function of time) to classify variable objects (e.g., \\citealt{palaversa13}). Because photometric pipelines can report false dimming events,  rare and deep dimming events have only been confirmed when the obscuration is periodic (e.g., the eclipsing disk OGLE-LMC-ECL-11893 with a 468 day period; \\citealt{dong14}) or multiple observation platforms observed the same drop in magnitude (e.g., RW Aur A;  \\citealt{rodriguez13}).  Here we focus on searching for potentially rare events and so we do not restrict our study to periodic stars or stars that are sufficiently variable that they can be found via statistical measurements of the magnitude distribution. % We carry out a multi-band search for obscuration events in a fairly large sample, over 40000 stars, the brighter stars in the 2MASS Cal-PSWDB. %2MASS Calibration Point Source Working Database. The Cal-PSWDB is uniformly observed and processed, and has  well characterized astrometric and photometric measurements.  It is also a multi-band photometric survey with all bands measured at each observation time. Variability is often studied in a single photometric wavelength band, however, if obscuration is due to extinction from dust, then the depth of a dimming event or eclipse can depend upon wavelength. For example it is possible to differentiate between star spots and obscuration using near-infrared  color \\citep{moralescalderon11,parks13,wolk13,cody14}. The eclipses of the OGLE-LMC-ECL-11893 and EE-Cep disk systems are deeper in $B$ band than $I$ band \\citep{galan12,dong14}. In a three band photometric study, it is also possible to check the uniformity or noise level of measurements at a particular time by comparing the measurements in the three bands. Here we carry out a search for dimming events in single bands and also search for dimming events that are red and so could be due to transient levels of extinction. ",
        "conclusions": "We have presented a search for dimming events in approximately 40,000 objects (mostly stars)  observed over a period of about 4 years in the 2MASS Calibration fields. We searched for day long dimming events that are significant compared to the standard deviation of the object's magnitude distribution. Our goal was to search for rare long dimming by an exotic object such as the occulting disk of OGLE-LMC-ECL-11893, consequently we did not restrict our search to objects with periodic variability. In a color independent search for day-long dimming events we found 46 previously unidentified eclipsing binaries %xxxx and 6 additional periodic variable stars.   We also found a few objects exhibiting dimming events but without strong color variations during dimming.   These could be eclipsing binaries but we have not observed enough eclipses to identify their period, or they could be something more exotic. We also found two active galaxies and a few objects displaying correlated color and magnitude variations suggesting that they are also extragalactic. Young stellar objects in the $\\rho$ Ophiucus field, previously studied by \\citet{parks13} also displayed dimming events, however no ordinarily non-variable or quiescent star in this region was detected with a dimming event. Two less obscured objects (Iso-Oph 151, a class III source and WL 13 a class II source) variable young stellar objects in this field displayed 0.2 mag deep, colorless two-day long dimming events. We failed to find any ordinarily quiescent object with a dimming event lasting more than a day. From the 4 year survey of approximately 40,000 objects (and excluding the young stellar objects) approximately 1 of 700 is an eclipsing binary, and approximately 1 of 4000 exhibits a dimming event but requires additional observations to classify.  We carried out a second search, looking for dimming events that are accompanied by reddening. In this search we did not restrict the search by the standard deviation of the object's magnitude distribution. This search primarily revealed two classes of objects, active galaxies that become bluer as they increase in brightness, and class II young stellar objects in the $\\rho$ Ophiucus region that exhibit large variations in brightness (up to a magnitude in $J$ band) that are also red ($\\Delta (J-K) \\sim 0.4$ mag). The young stellar objects found in both searches display color variations consistent with interstellar extinction, though in some cases there must be additional emission mechanisms influencing the brightness and color. Of the 40000 objects searched we found 14 active galaxies,  corresponding to approximately one of 3000 searched, exhibiting dimming events and correlations between brightness and color.  We did not find any dimming events greater than 0.2 magnitude deep and longer than a day outside the $\\rho$ Ophiucus region.    We did not find an object similar to OGLE-LMC-ECL-11893, ordinarily without variability, but exhibiting long, deep and red dimming events, in the 2MASS calibration fields (and including objects in the $\\rho$ Ophiucus field).  Because we failed to find any ordinarily quiescent object with a long (greater than a day) dimming event, we consider how we would eliminate more common objects in a future larger scale search. The red dimming event search effectively eliminates all but one eclipsing binary comprised of two stars with significantly different effective temperatures. The only targets that were found in both searches were the young stellar objects in the $\\rho$ Ophiucus star forming region, a single exceptional eclipsing binary and a rapidly varying object that might be an active galaxy. %The 2MASS Cal PSWDB contained a field covering part of the $\\rho$ Ophiucus star formation region. %A variability survey covering a larger fraction of the sky would be better at contrasting the fraction of  For objects found in our searches we computed correlation coefficients between color and magnitude and fit a slope to points in the color magnitude distribution.  We find that young stellar objects, active galaxies and objects displaying insignificant variations in color (and so are most likely eclipsing binaries) tend to be found in different regions of this plot, however these two quantities are not sufficient to definitely separate the populations. Most active galaxies can be eliminated with a limit on the dispersion in the magnitude distribution. Perhaps combining additional color information or by using higher moments of the magnitude distributions, or additional spectral information  the two populations could be better delineated and identified.  There is a long history of searching for and following up objects that temporarily brighten, such as novae or gravitational micro-lenses,  but there have been few systematic searches for dimming events. In this study we took care to discard multiple dimming events that occurred at the same time in the same calibration field. However, differential photometric methods, where the brightness of nearby stars are contrasted in the data reduction pipeline, can be used to minimize the number of spurious dimming events (e.g., see discussion by \\citealt{vaneyken11}). This would allow a more robust search for rare and single dimming events in a larger sample of stars and make it possible to develop strategies to follow up unidentified objects.  New and ongoing surveys are increasing the number of stars photometrically monitored, improving upon the cadence of observations and on the numbers of photometric bands for photometric monitoring. An object similar to  OGLE-LMC-ECL-11893 could be found in larger and  multi-color survey that searched for deep and red dimming events, a high cadence single band survey that examined the shape of the light curve during a single dimming event or computed the magnitude distribution in an eclipse \\citep{meng14}, or an extended (many year) lower cadence single band survey survey that examined the shape of the eclipse in the period-folded light curve  (\\citealt{dong14}; and how it was found). Rare events, such a transit of a planet hosting icy rings, like Saturn, or a disk with large dust grains (and a grayer extinction law) might not cause a red dimming event but could induce an atypically shaped eclipse.  Differentiating this from an eclipsing binary would require high cadence and high photometric precision (needed to detect the shape of the eclipse).  We note that our search criteria were not standard compared to those used in other surveys, so our search for variable stars in the 2MASS calibration fields is not complete.  For example, we did not follow up bright stars identified with large dispersions computed from the magnitude distribution in the light curve.  We searched here for rare dimming events using a multicolor near-infrared survey. An optical survey could survey a larger number of fainter stars, however a near-infrared survey would be particularly useful for studying star forming regions embedded in the Galactic plane.   Eclipsing disks associated with planet formation might better be found in a near-infrared survey that includes many young stars, whereas a search for ringed planets might be better done in the visible bands.  In our search of approximately 40000 light curves approximately 1 of 4000 objects (11 objects total) that have displayed dimming events remain unidentified.   As more common objects such as eclipsing binaries and active galaxies are classified in a continuously monitored field, it could become possible to discover more exotic objects displaying rare dimming events. Roughly scaling from our estimates, a search for dimming events in a billion good light curves would yield of order a million eclipsing binaries, a few hundred thousand active galaxies and a few hundred thousand unidentified objects (that would include young stellar objects). Amongst these could be a bright counterpart to OGLE-LMC-ECL-11893 or an analog of Saturn's rings.  \\subsection*{Acknowledgements} We thank Judy Pipher for helpful discussions. This work was in part supported by NASA grant NNX13AI27G."
    },
    "1402/1402.6983_arXiv.txt": {
        "abstract": "We investigate the bias and error in estimates of the cosmological parameter covariance matrix, due to sampling or modelling the data covariance matrix, for likelihood width and peak scatter estimators.  We show that  these estimators do not coincide unless the data covariance is exactly known. For sampled data covariances, with Gaussian distributed data and parameters, the parameter covariance matrix estimated from the width of the likelihood has a Wishart distribution,  from which we derive the mean and covariance. This mean is biased and we propose an unbiased estimator of the parameter covariance matrix. Comparing our analytic results  to a numerical Wishart sampler of the data covariance matrix we find excellent agreement. An accurate {\\it ansatz} for the mean parameter covariance for the peak scatter estimator is found,  and we fit  its covariance to our numerical analysis. The mean is again biased and we propose an unbiased estimator for the peak parameter covariance. For sampled data covariances the width estimator is more accurate than the peak scatter estimator. We  investigate modelling the data covariance, or equivalently data compression, and shown that the peak scatter estimator is less sensitive to biases in the model data covariance matrix than the width estimator, but requires independent realisations of the data to reduce the statistical error.  If the model bias on the peak estimator is sufficiently low this is promising, otherwise the sampled width estimator is preferable. ",
        "introduction": "The high precision required to probe the nature of dark energy, dark matter and modifications to gravity (e.g., Amendola et al. 2013) is driving cosmology to an era where the accuracy of parameter estimation will have to reach sub-percent levels. To meet this challenge  large-scale ground and space-based cosmological surveys are being planned and carried out which are optimised to deliver high statistical accuracy (e.g., VST-KiDS, DES, HSC, LSST, Euclid). For these surveys to be successful systematic biases will also have to be controlled to an unprecedented level, within the bounds set by the statistical uncertainty. The introduction of statistical uncertainty and systematic biases will have to be tracked at every step of the data analysis,  from observation to parameter estimation.    An aspect which has recently been receiving more attention in this process is  the final parameter estimation step when data, compressed into the form of power spectra or correlation functions, is compared with cosmological models and further compressed into estimates of the model parameters along with an estimate of their accuracy. In particular we need to have reliable, unbiased estimates of the  parameter covariance matrix, which is needed to demonstrate how accurate the parameters have been measured as well as delineating the volumes of parameter space where acceptable models reside. Beyond this, if we want to apply some form of model selection, for example investigating the Bayesian Evidence, we need to have an accurate representation of the posterior distribution of the parameters. If the parameter covariance matrix is biased by a poor estimator it will either over- or underestimate the actual errors and covariances of the measured parameters. In addition, a sub-optimal covariance estimator will itself have significant uncertainties which should be folded into the overall error budget.    There are two common  approaches to estimating the uncertainty on parameters derived  from cosmological data. One is to estimate the variance, or width, of the likelihood surface in parameter space. This can be done by mapping out the likelihood surface and numerically integrating on a grid,  or using a Monte-Carlo Markov Chain (e.g., Lewis \\& Bridle 2002) to sample the likelihood distribution and Monte-Carlo integrating.  A second approach is to generate   many independent  realisations of the survey, either by simulating or re-sampling the data, estimate the maximum likelihood parameter values for each realisation, and then use the scatter in the peak values as an estimate of the uncertainty in the result. In the limit of no bias or uncertainty in the model distribution, these should yield the same answer. However, as we shall show, if the likelihood distribution does not accurately model the distribution of the data both of these estimators will produce biased estimates of the parameter covariance matrix, and the width and peak scatter estimators will no longer coincide. In addition, biases in the likelihood function can significantly increase the uncertainty in the parameter covariance to the point that the error estimate is unreliable.   For data which follows a multivariate Gaussian distribution all of the statistical information is encoded within  the data covariance matrix and in particular its inverse, the precision matrix. If the data covariance is known exactly {\\em a priori}, it can be accurately inverted to find the precision matrix and the likelihood function is unbiased.  However, if the data covariance is not well known, or must be estimated, it will  be biased and this bias will propagate into the likelihood function through the precision matrix.    In Cosmology, where  the nonlinear evolution of density perturbations and effect of baryons, galaxy formation, stellar and AGN feedback affect the statistical properties of observables in a complex way (e.g., Semboloni et al. 2011, van Daalen et al. 2013), the data covariance matrix is usually  estimated by sampling independent realisations of the data, either from simulations or re-sampling the data using Jackknife (e.g. Tukey 1958) or Bootstrap (e.g., Efron 1979, Norberg et al. 2009) methods, or from models of the data covariance matrix which aim to contain nonlinearity and feedback effects  (e.g., Cooray \\& Hu 2001, Takada \\& Bridle 2007, Takada \\& Jain 2009, Hilbert et al. 2011, Kayo et al. 2012). Estimating the data covariance by sampling  independent realisations of the data will introduce a sampling variance which propagates into the precision matrix and parameter estimation. For Gaussian-distributed data the sample data covariance matrix follows a Wishart distributed (Wishart 1928), while the precision matrix follows an Inverse-Wishart distribution, first studied by Kaufmann (1967; see also, e.g., Taylor et al. 2013). Kaufmann derived the expectation value of the precision matrix and showed it was biased compared to the inverse of the expected data covariance, and that this bias diverges when the number of realisations approaches the size of the data set. Hartlap et al. (2007)  showed that this bias is found in cosmological precision matrices when the data covariance matrix is estimated from an ensemble of cosmological simulations, and suggested that the bias should be corrected for to avoid underestimating parameter errors.   Taylor et al. (2013) extended this analysis to estimating cosmological parameter covariance matrices and, assuming that the parameters were Gaussian-distributed in parameter space, showed to second order that the variance, or width, of the likelihood function would have an additional  uncertainty due to the sampled precision matrix. Dodelson \\& Schneider (2013)  further showed that if the parameter covariance was estimated from the scatter in the peak of the likelihood the parameter covariance would be biased high, to second order, due to the Inverse-Wishart scatter of the precision matrix. Likewise, Percival et al. (2014) showed a similar bias arises to second order in the width estimator.   In this paper we complete this analysis  by studying the distribution, bias and covariance of the parameter covariance matrix from the width and peak scatter estimators when the data covariance matrix is sampled from independent realisations of the data. In addition we investigate the effect of modelling the data covariance matrix,  where there is no sampling variance, but uncorrected biases will propagate into the parameter covariance matrix.  In Section 3 we discuss how to numerically  generate random realisations of the data covariance matrix to compare with our results. In Section 4 we derive the exact distribution for the parameter covariance matrix estimated from the width of the likelihood, and its bias and error. We also use our numerical results to propose an ansatz for the bias in the peak scatter estimator and a fit to its error. We study modelling of the data covariance and data compression in Section 5, and present our summary and conclusions in Section 6. We begin by reviewing estimators for the parameter covariance and how a bias in the data covariance matrix propagates. ",
        "conclusions": "\\begin{table*} \\centering \\caption{Summary  of new results for the statistics of parameter covariances presented in this paper. Equation numbers refer to results in this paper, following citations refer to previous second-order solutions.} \\begin{tabular}[t]{lll} \\hline\\hline & Width parameter covariance & Peak parameter covariance\\\\ \\hline Bias in the mean estimator & Eq. (\\ref{eq:DCC}) (2nd order Percival et al. 2013) &  Eq. (\\ref{eq:cov})  (2nd order Dodelson \\& Schneider 2013)\\\\ Variance of biased estimators &  Eq. (\\ref{eq:widthcov})  (2nd order Taylor et al. 2013) &     \\\\ Variance of unbiased estimators & Eq. (\\ref{eq:widthcovunbiased}) & Eq. (\\ref{eq:peakcov})\\\\ \\hline \\end{tabular} \\label{tab:structure} \\end{table*}  As Cosmology progresses into the age of high precision measurements, probing the nature of dark energy and dark matter, and looking for evidence of modified gravity,  accurate estimates of cosmological parameter errors and covariances require accurate data covariance matrices. Due to the nonlinearity in cosmological fields, and the complexity and nonlinearity of data  analysis pipelines, a common approach is to estimate the data covariance from an ensemble of realisations of the survey, usually using simulations or empirically from the data from Jackknife or Bootstrap resampling. In maximum likelihood parameter estimation, errors and covariances  are derived from the width of the likelihood surface or from the scatter in the peak of the likelihood from a set of realisations. However,  statistical uncertainty and bias in the estimation of the data covariance matrix will propagate into additional errors on the measurement of cosmological parameters.  In this paper we have found expressions for the change in the parameter covariance for both width and peak scatter estimators, assuming that the data and parameters are Gaussian-distributed,  due to a change in the data covariance matrix from its expectation value. In particular, the width error estimate is biased at first order while the peak scatter error estimate has a second order bias, and so the error estimates from both estimators will not in general coincide.   If the data covariance matrix  for $N_D$ data points is estimated from $N_S$ independent random Gaussian-distributed realisations of the data, it will be Wishart-distributed with $N_S-1$ degrees of freedom. We have shown here that the parameter covariance matrix for $N_P$ parameters,  estimated from the width of the likelihood, is also Wishart-distributed with $N_S-N_D+N_P-2$ degrees of freedom. With the full distribution of the parameter covariance matrix we have derived its  expectation value (equation \\ref{eq:DCC}) and its covariance (equation \\ref{eq:widthcov}). In general, the estimated parameter covariance will be larger than the optimal parameter covariance matrix, estimated from an infinite number of realisations, but we can renormalise the estimated parameter covariance matrix to find an unbiased estimate. The fractional variance of the unbiased width estimator then scales as $\\nu^2 = 2/(N_S-N_D+N_P-1)$. Using our expressions for the change in the parameter covariance due to the change in the data covariance, and a Wishart sampler to generate realisations of the data covariance matrix for $N_S$ realisations, we have numerically estimated the expectation and variance of the parameter covariance and found excellent agreement with our analytic results. Table \\ref{tab:structure} summarises the key expressions  for biases and variances of the parameter covariance matrices.   The distribution of the peak scatter parameter covariance matrix, given a Wishart-distributed data covariance, does not appear tractable but we have found a very accurate ansatz for its expectation value (equation \\ref{eq:cov}). Again we have used this expectation value, which is again larger than the optimal estimate, to renormalise and find an unbiased peak scatter estimator. The covariance of the parameter covariance matrix is more complex, as we need to have independent realisations of the data to estimate both the data covariance matrix and to reduce the uncertainty in the estimate arising from the random scatter in the position of the likelihood peak. For a fixed number of realisations, we find the optimal split between the data covariance and peak scatter estimates and fit an accurate expression for the covariance of the unbiased estimate (equation \\ref{eq:peakcov}). We have found that the parameter covariance estimated from the peak scatter is generally larger than the width estimate. For the unbiased estimators the variance of the optimal peak scatter estimator is also larger than that of the width estimator, and so in general we advocate the width estimator over the peak scatter estimator.     The Wishart bias for both peak and width estimators vanishes if the number of data points equals the number of parameters, however the  scatter on both  scales as $\\sqrt{2/N_S-1}$. We can achieve $N_D=N_P$ by data compression, since expanding the parameter space does not help. In our analysis data compression is equivalent to assuming a model for the data covariance matrix, and the accuracy in the parameter covariance in both cases  is determined by the accuracy of the model. The width parameter covariance estimator is more sensitive to biases in the data covariance model than the second order peak scatter  estimator to the assumed data covariance.This implies that for data compression, or modelling, the model data covariance matrix will be more accurate for the peak estimator than the width estimator The peak scatter estimator requires independent realisations to estimate the parameter covariance and we need  $N_S > 1+2/(\\nu^2-b_P^2)  $ realisations to reach a fractional accuracy of $\\nu$ in the parameter covariance matrix, where $b_P$ is the fractional  bias from the model data covariance matrix.   In current and future cosmological surveys where the number of data points, for example power spectra bandpass modes or correlation function points in redshift bins, will grow from hundreds to thousands or tens of thousands,  to control the accuracy of errors we will have to generate $N_S > N_D-N_P+2/\\nu^2+1$ independent realisations of the data, or accurately model the data covariance matrix. If we sample independent realisations to estimate the data covariance, we have shown that we can remove the inherent bias in estimates of the parameter covariance matrix and control its error, preferring a width estimator over a peak scatter estimator. If we model the data covariance matrix, or apply data compression, we find the lower sensitivity of the peak scatter estimator makes it preferable to the width estimator, but the peak scatter may need large numbers of  independent realisations to reduce the random error on the parameter covariance if the bias is significant. Further work is required to assess if the model bias is low enough to apply the peak scatter estimator. If not, the sampled unbiased width estimator is preferred."
    },
    "1402/1402.3568_arXiv.txt": {
        "abstract": "We present the results of near-infrared spectroscopy of \\ha\\ emitters (HAEs) associated with two protoclusters around radio galaxies (PKS1138-262 at $z$=2.2 and USS1558-003 at $z$=2.5) with Multi-Object Infrared Camera and Spectrograph (MOIRCS) on the Subaru telescope. Among the HAE candidates constructed from our narrow-band imaging, we have confirmed membership of 27 and 36 HAEs for the respective protoclusters, with a success rate of 70 per cent of our observed targets. The large number of spectroscopically confirmed members per cluster has enabled us for the first time to reveal the detailed kinematical structures of the protoclusters at $z$$>$2. The clusters show prominent substructures such as clumps, filaments and velocity gradients, suggesting that they are still in the midst of rapid construction to grow to rich clusters at later times. We also estimate dynamical masses of the clusters and substructures assuming their local virialization. The inferred masses ($\\sim$10$^{14}$M$_\\odot$) of the protocluster cores are consistent with being typical progenitors of the present-day most massive class of galaxy clusters ($\\sim$10$^{15}$M$_\\odot$) if we take into account the typical mass growth history of clusters. We then calculated the integrated star formation rates of the protocluster cores normalized by the dynamical masses, and compare these with lower redshift descendants. We see a marked increase of star-forming activities in the cluster cores, by almost three orders of magnitude, as we go back in time to 11 billion years ago; this scales as (1$+$$z$)$^6$. ",
        "introduction": "In protoclusters at $z$$>$2, characteristic relations seen in low-$z$ clusters ($z$$<$1) such as the colour-magnitude relation, break down as galaxies enter into their formation phase \\citep{Kodama:2007}. Since the galaxies in those protoclusters are destined to evolve into early-type galaxies in rich clusters today, the protoclusters provide us with unique laboratories in which to investigate directly the formation mechanisms of early-type galaxies and their environmental dependence, through comparison with field galaxies at similar redshifts. It is therefore essential to investigate characteristics of protoclusters at $z$$>$2 systematically, in order to know how the star forming (SF) activities in high density regions at high $z$ are intrinsically biased, and how they are affected externally by their surrounding environments to establish the strong environmental dependence seen in the present-day Universe.  With this motivations, we have been conducting a systematic study of protoclusters at $z$$>$1.5 with Subaru, as the project `Mahalo-Subaru' ({\\it Mapping HAlpha and Lines of Oxygen with Subaru}; for more detail see \\citealt{Kodama:2013}). We have conducted narrow-band (NB) imaging with many customized NB filters, and have successfully identified \\ha\\ or \\oii\\ emitters candidates that are physically associated with the protoclusters. The following two objects are among the richest systems ever identified: USS1558--003 at $z$$=$2.53 (\\citealt{Hayashi:2012}, hereafter H12), and PKS1138--262 at $z$$=$2.16 (\\citealt{Koyama:2013}, hereafter K13). Since they both show large excesses in number densities of SF galaxies, these protoclusters are probably still in the vigorous formation process. Our observations have revealed high SF activities towards the cores of protoclusters at $z$$>$2. The peak of SF activity traced by the line emitters is shifted from dense cluster cores to lower density outskirts and filamentary outer structures with time from $z$$\\sim$2.5 to $z$$\\sim$0.4, indicating the inside-out growth of clusters.  These results are all intriguing but we need to confirm them and investigate the physical properties of protocluster galaxies in much greater detail with spectroscopic follow-up observations. This Letter reports the first results of our near-infrared (NIR) spectroscopies of HAEs in the two richest protoclusters at $z$$>$2. We first describe our targets and the method of spectroscopic observations and then present the kinematical structures of spectroscopically confirmed members. We also discuss SF activities in these two systems through comparison with lower redshift counterparts. We assume a $\\Lambda$-dominated cosmology with of $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7 and $H_0$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Kinematical structures of distant protoclusters provide essential information on the mass assembly history of galaxy clusters. Fig. \\ref{fig2} presents the spatial and redshift (or radial velocity) distributions of the HAEs (and 15 confirmed LAEs in the case of PKS1138). The blue and red symbols separate the members according to their redshifts or radial velocities: those approaching or receding, respectively, with respect to the RGs.  From the velocity dispersion, we now estimate R$_{200}$ of each protocluster, which is the radius within which the averaged matter density is 200 times larger than the critical density. The dynamical mass of each system (M$_{cl}$) is also measured using the virial theorem \\citep{Finn:2005}, assuming local virialization.  \\subsection{PKS 1138-262 ($\\mathbf{z}$=2.16)} PKS1138 is among the most familiar protocluster at $z$$\\sim$2, and has been intensively studied by many researchers (e.g. \\citealt{Kurk:2000}).  We see some velocity structures across the protocluster (Fig. \\ref{fig2}a) as well as the spatial structure. The galaxies near the RG within the dashed circle denoted as C1 seem to be relatively well mixed in velocity. However, the outer region are more structured. The most prominent structure is a linear filament extending towards south-east direction from the RG. This almost perfectly aligned filament is dominated by approaching galaxies with respect to the RG, suggesting that those galaxies are falling onto the protocluster core along the filament and penetrating into the very centre. On the other hand, the north-west and east areas are preferentially occupied by receding populations. The south-west complex further away from the protocluster centre near the bottom right corner of the figure consists solely of approaching galaxies, suggesting the existence of a coherently moving group or filament.  All these spatial and kinematic structures seem to suggest that the inner part of the cluster centred on the RG ($<$0.5Mpc) is already collapsed and nearly virialized, while the outer regions are still highly structured and are at the early phase of assembly towards the protocluster core. The radio galaxy PKS1138 is right at the junction of the infalling filaments at the dynamical centre.  Upon the assumption of virialization in the core, the dynamical mass of the core is estimated to 1.71$\\times$10$^{14}$ M$_\\odot$ from the velocity dispersion of 683 km s$^{-1}$ within C1. The X-ray observation of this system with the High Resolution Imager on {\\it ROSAT} show that the emitted energy is 6.7$\\pm$1.3$\\times$10$^{44}$ ergs s$^{-1}$ in 2--10 keV band corresponding to the dynamical mass of $\\sim$10$^{14}$ M$_\\odot$. It is consistent with our result, although the X-ray emission is contaminated by active galactic nuclei \\citep{Carilli:1998, Pentericci:2002}.  \\subsection{USS1558-003 ($\\mathbf{z}$=2.53)}  USS1558 is the densest protocluster ever known to date at $z$$>$2, and was first discovered by \\citet{Kajisawa:2006} and H12. H12 identified three groups of HAEs aligned along the north-east--south-west directions: a loose group around the RG, the richest clump at 3.5 arcmin away from the RG, and a small clump in between the two. This work revises that previous grouping to two clumps C1 and C2 as shown in Fig.\\ \\ref{fig2}(b) because of the poor kinematical separation between the latter two clumps, identified only spatially in the previous study. We have now confirmed spectroscopically that the two groups are actually located at the redshift of the RG and are embedded in the large-scale structure. C2 is particularly rich, as it confines 19 spectroscopically confirmed HAEs and 12 more candidates (not confirmed yet) within radius 0.6 Mpc and is the densest system ever identified at high redshifts ($z$$>$2). We note that the RG itself seems to be offset from this densest clump, unlike the PKS1138 cluster.  Judging from the very high densities of the HAEs in compact areas, we could reasonably assume a local virialization in these regions, and the corresponding dynamical masses and R$_{200}$ are estimated and listed in Table \\ref{tab2}. The dynamical mass of C2 is estimated to 0.87$\\times$10$^{14}$ M$_\\odot$ from its velocity dispersion of 574 km s$^{-1}$.  We find that there is a large-scale velocity gradient across the cluster in the direction of the group alignment (north-east--south-west). The velocity distribution and central value of the south-west group (C2) are blue-shifted from those of the north-east group (C1) that hosts the RG (Fig.\\ \\ref{fig2}; see also Table \\ref{tab2}). Therefore those groups are probably physically aligned and gravitationally pulling each other closer. They would eventually merge together and become a single rich cluster in the near future.  \\begin{figure} \\centering \\includegraphics[width=80mm]{fig3.eps} \\caption{(a) Integrated SFRs in the cluster cores, $\\Sigma$SFR, normalized with the cluster dynamical masses (upper panel) and (b) cluster dynamical masses (lower panel) are plotted as a function of redshift. The open diamonds show the measurements for the Mahalo-Subaru cluster sample including this work, calculated within $R_{200}$ of each cluster \\citep{Koyama:2010, Koyama:2011, Hayashi:2011, Tadaki:2012}. The filled diamonds indicate the values within 0.5$R_{200}$ to match the definition with other previous measurements for a direct comparison. The previous works for eight clusters at $z$=0.1--0.9 are shown by squares (compilation by \\citealt{Finn:2005}). The grey and open squares separate those clusters according to their dynamical masses, as shown in the lower panel. Note that SFRs are measured from \\ha\\ line strengths for all the clusters except for the $z$=1.46 and $z$=1.62 clusters which are based on \\oii\\ lines. Note also that $R_{200}$ is not fully covered for the $z$=0.81 cluster. The cluster mass of J0218 cluster at $z$=1.62 is adopted from \\citet{Tanaka:2010}. The red dotted curve shows the relation as a function of redshift, scaling as (1$+$$z$)$^6$. In the lower panel, the red line and the pink zone show the typical mass growth history of massive cluster haloes with 1--2$\\times$10$^{15}$\\msun\\ predicted by theoretical models \\citep{Shimizu:2012, Chiang:2013}. The pink zone corresponds to $\\pm$1$\\sigma$ scatter around the median values. } \\label{fig3} \\end{figure}  \\subsection{Cosmic evolution of $\\mathbf{\\Sigma}$SFR/M$_{\\rm\\mathbf{cl}}$}  Finally, we investigate the cosmic SF history in galaxy clusters represented by the integrated SFR ($\\Sigma$SFR) normalized by cluster dynamical mass (Fig. \\ref{fig3}a). In order to compare our results directly with the previous works compiled by \\citep{Finn:2005}, we sum up individual SFRs of the \\ha\\ (or \\oii) emitters within 0.5$R_{200}$, including the candidates whose membership has not been confirmed yet. In this work, we calculate $\\Sigma$SFR and M$_{cl}$ for the main body of the PKS1138 cluster (C1) and the richest clump of USS1558 cluster (C2), respectively. We assume an uniform dust extinction of A$_{H\\alpha}$=1 and A$_{\\rm\\,[O{\\sc II}]}$=1.76A$_{H\\alpha}$. To evaluate the uncertainties from this assumption, we also estimate $\\Sigma$SFR within R$_{200}$, employing the mass-dependent correction for dust extinction \\citep{Garn:2010}. Although the absolute values of $\\Sigma$SFR may have large systematic errors, due to various factors such as sampling bias, active galactic nucleus contribution and so on, the relative differences among different clusters that we see here are more reliable. \\citet{Koyama:2010, Koyama:2011} found that $\\Sigma$ SFR/M$_{cl}$ in cluster cores increases dramatically to $z$$\\sim$1.5 and scales as $\\sim$(1+$z$)$^6$ (see also \\citealt{Smail:2014}). We find that this trend extends to even higher redshifts to $z$$\\sim$2.5, as our values of $\\Sigma$SFR/M$_{cl}$ within $R_{200}$ estimated in this work seem to more or less follow the extrapolated curve of the redshift evolution.  In such comparison of clusters at different redshifts, we must be sure that we are comparing the right ancestors with the right descendants. In fact, clusters grow in mass with cosmic times by a large factor and therefore we should compare galaxy clusters taking into account such mass growth. Figure \\ref{fig3}(b) show cluster masses as a function of redshift. The red line and the pink zone show the mass growth history of massive cluster haloes predicted by cosmological simulations \\citep{Shimizu:2012, Chiang:2013}. The data points show the measurements of dynamical masses of real clusters used for comparison. It turns out that our protoclusters at $z$$>$2 have large enough masses to be consistent with the progenitors of the most massive class of clusters, like Coma. The lower-$z$ clusters, shown with filled squares also follow the same mass growth curve. Therefore we argue that we are comparing the right ancestors with right descendants and the redshift variation of the mass-normalized SFRs seen in the upper panel can be seen as the intrinsic cosmic SF history of the most massive class of clusters.  In this Letter, we have presented the kinematical structures of the two richest protoclusters at $z$$>$2, and extended the cosmic evolution of $\\Sigma$SFR/M$_{cl}$ back to $z$$>$2 or 11 Gyrs ago, based on the intensive multi-object NIR spectroscopy of the NB-selected SF galaxies. In our forthcoming Paper I\\hspace{-.1em}I (in preparation), we will discuss the physical properties of these galaxies (such as gaseous metallicities, ionizing states and dust extinctions) using the multi-emission-line diagnostics, and compare them with those in the general field at similar redshifts."
    },
    "1402/1402.2758_arXiv.txt": {
        "abstract": "{We present the detailed modelling of line spectra emitted from galaxies at redshifts 0.2$\\leq$ z$\\leq$ 2.3. The spectra account only for  a few oxygen  to \\Hb line ratios. The results show that   [OII]3727+3729/\\Hb and [OIII]5007+4959/\\Hb are not sufficient to  constrain the models. The data  at least of  an auroral  line, e.g. [OIII]4363, should be known. We have found  by modelling the  spectra  observed from ultrastrong emission line galaxy and faint galaxy  samples, O/H relative abundances  ranging between 1.8$\\times$10$^{-4}$ and 6.6$\\times$10$^{-4}$.}   \\keywords {radiation mechanisms: general --- shock waves --- ISM: O/H abundances ---  galaxies: starburst --- galaxies: high redshift}   \\titlerunning{Modelling galaxies at 0.2$<$z$<$2.3} \\authorrunning{M. Contini} ",
        "introduction": "Observations  of line spectra emitted from galaxies at relatively high redshifts are now available (see Contini 2013b, hereafter Paper I, and references therein). Metallicities, in terms of the relative abundance of the  heavy elements to H, have been  calculated by a detailed modelling of the line ratios, providing  some information about galaxy evolution, star formation rates,  luminosities etc.  The spectra are relatively poor in number of lines at high z because only the strongest  ones are observable. The modelling procedure, therefore, becomes problematic  regarding model degeneracy.  Line ratios corresponding to different elements in a spectrum depend  on  the physical parameters and on the  relative  abundances of the elements. The O/H ratio shows generally the highest relative abundance  compared with that of the other  heavy elements  whose lines are observed in the UV-optical -IR frequency range. Different solar O/H  are reported by Asplund et al (2009), Allen (1976) and Anders \\& Grevesse (1989), namely, O/H=4.9 10$^{-4}$,  6.6 10$^{-4}$ and 8.5 10$^{-4}$, respectively. Table 1 shows that Allen (1976)  relative abundance  values  range  between the two more recent results. Therefore, we will refer to Allen (1976). Yet, the relative abundances are calculated consistently for each spectrum. So the  values which appear in Table 1 are important only as references for  discussions.   \\begin{table} \\caption{The solar element abundances} \\begin{tabular}{ccccccc} \\hline  \\hline \\ element & Allen & Anders \\& Grevesse & Asplund et al.  \\\\ \\         & (1976)       & (1989)      & (2009)      \\\\ \\hline \\    H    &  12          &    12       & 12           \\\\ \\  C      &  8.52        &  8.56       &  8.43       \\\\ \\  N      & 7.96         &  8.05       &  7.83       \\\\ \\  O      & 8.82         &  8.93       &  8.69  \\\\ \\   Ne    & 8.            & 8.09       & 7.93 \\\\ \\ Mg      & 7.4          &  7.58       &  7.6 \\\\ \\ Si      & 7.52         &  7.55       & 7.51   \\\\ \\ S       & 7.2          &  7.21       & 7.12   \\\\ \\ Cl      &  5.6         &   5.5       & 5.5   \\\\ \\ Ar      & 6.52         &   6.56      & 6.4   \\\\ \\ Fe      & 7.5          &   7.67      & 7.5  \\\\ \\hline\\\\ \\end{tabular} \\end{table}  Line ratios  from the same element in different ionization stages e.g. [OIII]5007+/[OII]3727+ (the + indicates that the $\\lambda$$\\lambda$5007,4959 and $\\lambda$$\\lambda$3727,3729 doublets are summed up), constrain the physical parameters such as the photoionization flux reaching the gas, the temperature of the line emitting gas, etc . The oxygen line ratios to \\Hb   constrain  the O/H  relative abundance. So the  \"direct\" or \\Te ~ method (Seaton  1975, Pagel et al. 1992, etc) is  used to obtain  O/H from the observed oxygen  to \\Hb line ratios. By this method, the ranges of the  gas physical  conditions are   chosen among those   most  suitable to the observed line ratios, e.g. the  temperature   is calculated from  [OIII]5007+/[OIII] 4363, the  temperature and the density  ranges are  constrained by  [OII]3727+/\\Hb considering, in particular,  that the critical density for collisional deexcitation of [OII] is $<$ 3000 \\cm3.  The density is constrained also by [SII] $\\lambda$6717/$\\lambda$6731 line ratio, when observed.  In  Paper I we have modelled the spectra of galaxies   at redshifts between 0.001 and 3.4, on the basis of   at least hydrogen, oxygen and nitrogen lines. We have adopted a composed model which accounts for the photoionization  and for shocks.  In the present  paper we  deal with spectra showing   a few oxygen  line ratios to \\Hb, i.e. the [OIII] 5007+4959/\\Hb, [OIII]4363/\\Hb and [OII]3727+/\\Hb observation data reported by Kakazu et al (2007) for a sample of  {\\it  ultrastrong emission line galaxies} at z$\\sim$1   and the [OIII] 5007+4959/\\Hb and [OII]3727+/\\Hb observation data reported by Xia et al. (2012) for a sample of {\\it faint galaxies} at 0.6$\\leq$ z $\\leq$ 2.4. Moreover, we have added a subsample of the stacking spectra of emission line galaxies at 1.3$\\leq$ z$\\leq$2.3 reported by Henry et al (2013). The stacking method is justified by the increasing number of  data observed in  the different galaxy samples.  Our aim is to discuss modelling degeneracy. We would like to find out the  smallest set of oxygen line ratios to \\Hb suitable to constrain the O/H relative abundance. Moreover, we will compare the O/H results  calculated by  detailed modelling  with those obtained by the other methods, e.g. the direct method  and the metallicity calibrators (Perez-Montero \\& Diaz 2005 and references therein).  We will model the spectra, even those showing a few lines, by the method adopted for  the spectra rich in number of lines (Paper I and references therein). Namely, the gas is  ionized and heated by the black body radiation flux from the stars in the case of a starburst (SB) galaxy or by a power-law radiation flux from  the active nucleus in active galactic nuclei (AGN).  It  was found that galaxies at a relatively high z are the product of merging, therefore a shock dominated regime is assumed leading to compression and heating of the emitting gas downstream of the shock front. Collisional ionization and heating prevail on the radiation processes  at relatively high shock velocities.  When the observed line spectrum of a galaxy contains a few hydrogen lines, oxygen lines from no more than two ionization stages, while  the lines from the other elements are missing, the results of the calculation process  can confront degeneracy. This refers at least  to the abundances of the elements that are relatively strong coolants, whose lines are not seen. Carbon lines are  not available in the optical range. Neon is hardly included into dust grains and molecules due to its atomic structure, therefore neon could be useful to  investigate the evolution of the heavy elements with z, but its lines are weak. Even if Ne/H $\\geq$ N/H,  we will consider nitrogen as the second  important heavy element because the N  lines (e.g. [NII] 6584, 6548) observed from luminous galaxies at redshifts as high as  z$\\leq$3.5 (Paper I) are relatively strong.  We will check this issue  by the detailed modelling  of the  Kakazu et al (2007) sample  of ultrastrong line emission galaxies and the Xia et al. (2012) faint galaxy sample.  In this paper we will first model the spectra adopting a solar N/H  and discuss degeneracy by reducing the N/H relative abundance. In fact, the results obtained in Paper I dealing  with spectra rich enough in number of lines,  show that N/H splits from $\\geq$ 10$^{-4}$ to $<$ 10$^{-5}$  at  redshifts in the 0.2$\\leq$ z $\\leq$ 1 range.  The  observed spectra  are relatively  poor because  many significant lines are missing (e.g. [NeIII], HeII, [OI], [NII] and [SII] etc) . Moreover,  the trend of  [OIII]5007+/\\Hb versus [OII]3727+/\\Hb  (Fig. 1) indicates that the distribution of the data follows the ionization parameter rather than the O/H  relative abundance, with some scattering due to the different physical conditions in the different objects.  To constrain the models by a first choice of the  physical  parameters, we have used the grids (Contini \\& Viegas 2001a,b) calculated previously by the  SUMA code for SBs and AGNs. Then, we have refined the models in order to reproduce the observed line ratios.  In Sect. 2  the modelling method is presented. In Sect. 3 the Kakazu et al. (2007) sample is modelled. In Sect. 4 we  deal with the Xia et al. (2012)  spectra. The results of Henry et al (2013) sample galaxies   appear in Sect. 5. Discussion and concluding remarks follow in Sect. 6.  \\begin{figure} \\includegraphics[width=9.0cm]{fig1.eps} \\caption{The observed [OIII]5007+/\\Hb versus [OII]3727+/\\Hb by Kakazu et al : black asterisks; by Xia et al : blue circles: Henry et al. : green asterisks. Red squares indicate the observed  [OIII]4363/\\Hb versus [OII] 3727+/\\Hb line ratios by Kakazu et al.} \\end{figure}   \\begin{figure} \\includegraphics[width=8.8cm]{fig2.eps} \\caption{Top panels : the profile of the electron temperature (solid line) and of electron density (dashed line); bottom panels : the profile of  the  fractional oxygen ions : O$^0$/O (black solid), O$^+$/O (red dashed), O$^{2+}$/O (blue dot-dashed), O$^{3+}$/O (green solid), O$^{4+}$/O (magenta solid) and of H$^+$/H (red solid) throughout a cloud corresponding to Kakazu et al ID 270 galaxy. The shock front is on the left of the left panel; the photoionization flux reaches the right edge of the right panel. The model corresponds to O/H=6.5 10$^{-4}$. } \\includegraphics[width=8.8cm]{fig3.eps} \\caption{The same as for Fig. 2 for models calculated by  O/H=1.8 10$^{-4}$. } \\end{figure} ",
        "conclusions": "The spectra observed from the  galaxies presented    by Kakazu et al. and by Xia et al.  surveys are  employed to calculate the physical conditions of the emitting gas by a detailed modelling  of ultrastrong emitting line galaxies and faint galaxies, respectively, at intermediate redshifts. Moreover, we have modelled the spectra presented by Henry et al which were obtained from the stacking of 85 luminous galaxies.  The results of the most significant parameters are shown in Fig. 4. The O/H ratios are shown in units of 10$^{-4}$ and \\Te in units of 10$^4$ K. All the results presented in Tables 2,  4 and 5  are compared with those calculated for a much larger sample of galaxies in Fig. 5.   \\begin{figure*} \\begin{center} \\includegraphics[width=7.8cm]{fig5a.eps} \\includegraphics[width=7.8cm]{fig5b.eps} \\includegraphics[width=7.8cm]{fig5c.eps} \\includegraphics[width=7.8cm]{fig5d.eps} \\includegraphics[width=7.8cm]{fig5e.eps} \\includegraphics[width=7.8cm]{fig5f.eps} \\caption{The evolution of different parameters with z (adapted from Paper I, Fig. 3). Top left : \\n0 and \\Vs; top right : the radiation parameters $F$ (in units of 10$^{10}$ photons cm$^{-2}$ s$^{-1}$ eV$^{-1}$ at the Lyman limit), the ionization parameter $U$ and the temperature of the stars (in units of 10$^4$ K); middle left : the relative abundances in units of 10$^{-4}$; middle right : \\Haa and $D$; bottom  left: the radius R (in kpc) of the NLR  in AGN  and of the emission regions in SBs and the \\Ha luminosity (L$_{H\\alpha}$): bottom right : the metallicity versus L$_{H\\alpha}$. Symbols : green + , green x refer and green asterisks refer to the Kakazu et al,  Xia et al and Henry et al samples, respectively. Red circles represent SB galaxies (Viegas et al. 1999, Contini 2013a,  Ramos Almeida et al 2013, Capetti et al 2013, Winter et al. 2010); red circles encircling a x refer to the SBs in the optically faint ULIRGs (Brand et al. 2007); red  triangles : SB  galaxies in the LINER sample (Contini 1997); red dots :  HII regions in star forming galaxies galaxies (Kobulnicky \\& Zaritsky 1999); red hexagram : ULIRG in QSO 2222-0964 (Krogager et al. 2013); black asterisks  : the AGNs (Ramos Almeida et al. 2013, Contini 2013a, Schirmer et al. 2013, Winter et al. 2010); black squares refer to the optically faint ULIRGs (Brand et al. 2007); blue filled triangles : the AGNs belonging to the LINER sample (Contini 1997); black pentagram  : ULIRGs (Krogager et al 2013). For Ne/H. Red plus : starbursts; blue cross : LINER AGN ; black dots : AGN. }  \\end{center} \\end{figure*}    We have investigated whether the  line ratio modelling  results  are constrained by the number and type of the observed lines, in terms of degeneracy. We have discussed, in particular, the  Xia et al sample which do not show line ratios other than  [OIII]5007+/\\Hb and [OII]3727+/\\Hb. The models  are  hardly constrained. To  investigate the results found in Sect. 3  about the  leading   role of the [OIII] 4363 /\\Hb line ratio, we presented the following test. A first  grid of models  was run adopting O/H ratios  as close to solar (Allen 1976) as  those adopted  to fit the sample of galaxies  which appear  in Paper I and in the  Kakazu et al. sample. The O/H relative abundances were readjusted in order to best fit the data  in tune with the other input parameters. The results appear in Table 4. Then, we recalculated the spectrum of the Xia et al ID 195 galaxy  by O/H= 10$^{-4}$ in agreement with the results obtained by the  methods adopted by the other authors. The observed [OIII]5007+/\\Hb and [OII] 3727+/\\Hb  ratios were reproduced by both  high and low O/H models. To remove degeneracy we invoked the observed weakness of the [OIII] 4363 line. We  conclude that spectra referring only to  [OIII]5007+/\\Hb and [OII] 3727+/\\Hb cannot be modelled with precision  without any further indication about e.g. the FWHM of the line profiles or the intensity of  some other lines.  \\subsection{Results for different N/H}  Let us now investigate the role of the abundances of elements  different from oxygen. Actually, the Kakazu et al and Xia et al samples were chosen by our investigation because, due to the observing difficulties at those z, the lines corresponding to heavy elements, other than oxygen,  were not reported. Those lines are generally weaker than the  oxygen ones, and/or they cannot be easily deblended. Nevertheless, the gas is generally  composed by the most prominent elements. We have used in the  present models  Allen (1976) solar abundances (He/H=0.1, C/H=3.3 10$^{-4}$, N/H=9.1 10$^{-5}$, O/H=6.6 10$^{-4}$, Ne/H=1. 10$^{-4}$, Mg/H=2.6 10$^{-5}$, Si/H=3.3 10$^{-5}$, S/H=1.6 10$^{-5}$, Cl=4. 10$^{-7}$, A/H=3.3 10$^{-6}$, Fe/H=3.2 10$^{-5}$).  In Fig. 5 we  have added to the modelling results calculated in Paper I   the results  calculated for the galaxy samples reported in the present work. The  relatively large galaxy sample presented in Paper I,  contains galaxies which were selected  among those including at least the [NII]6548+6584 doublet, in order to minimize degeneracy. The O/H relative abundances presented in Tables 2, 4 and 5 follow the trend found for different types of galaxies (Fig. 5, left middle diagram).  Changing the  abundance of one of the heavy elements relative to H, in particular for strong coolants   e.g. O, N, etc, the results of  the emitted line ratios would  change. In fact, the relative abundance of each element affects not only the line intensity but also the cooling rate of the gas in the recombination zone.  Let us investigate whether our models which adopt solar abundances for the elements corresponding to unobserved  lines  lead to trustful results. This is a critical test which could invalidate  most of  the results  presented in this paper. In Paper I we  have obtained by modelling the spectra reliable N/H and O/H for each galaxy. The N/H versus O/H relative abundances are shown in Fig. 6. Even with a large scattering,  the data show that there is an increasing trend of N/H with O/H. So in order to investigate  degeneracy  which   may  result by  changing  the  N/H input in a spectrum, we have run models with N/H lower than solar  for  galaxies which show a relatively low O/H.  We report in Table 6 the oxygen line ratios observed and calculated for two galaxies : ID 9 (NB912)  from the Kakazu et al sample and ID 339 from the Xia et al sample. Both refer to a relatively low O/H (Tables 2 and 4). Table 6 shows that the N/H  relative abundance does not affect the results  as much as to imply   a new set of the other  input parameters. The observational error  is $<$ 10\\% for  most of the lines. On the other hand, changing the parameters such as \\n0,  \\Vs, and/or $F$ can lead to strong differences in the (oxygen) line ratios (see Contini \\& Viegas 2001a,b). Therefore,  the results of the present work  are  safe concerning the physical conditions and the O/H relative abundance.  As  confirmed by  Fig. 6, different N/H  can correspond to the solar O/H for many galaxies. The  different  N and O trends   can be understood following Edmunds \\& Pagel (1978), namely, by the distinction between 'primary' elements such as oxygen and the   'secondary' ones. Nitrogen is a secondary element contaminated, however, by a non negligible primary component. Edmunds \\& Pagel  suggest that {\\it although  oxygen is instantaneously recycled by the supernova synthesis, nitrogen could be released by longer-lived stars and hence will appear in the ISM with a delay.}  In other words, the scattering of the N/O abundances at a relatively constant O/H which appears at z$\\leq$ 0.1 can be explained by  a  retarded release of N  produced in low mass longer-lived stars compared with O produced in massive, short-lived stars (Mouhcine \\& Contini 2002).    \\begin{table} \\centering \\caption{Comparison of line ratios to \\Hb=1 calculated  with  different N/H} \\tiny{ \\begin{tabular}{lccccccccc} \\hline  \\hline ID     &   z  & [OIII]5007+  &  [OIII]4363 & [OII]3727+ &  O/H &       N/H  \\\\ &      &              &             &            & 10$^{-4}$ & 10$^{-5}$ \\\\ \\hline 9 $^1$ & 0.833& 5.97         & $<$0.12     & 1.57       & -    &  -   \\\\ m1     &      & 6.0          & 0.037       & 1.62        & 2.&10.  \\\\ m2     &      & 6.1          & 0.038       & 1.63       & 2.&6.  \\\\ m3     &      & 6.16         & 0.039       & 1.65       & 2.&2.  \\\\ m4     &      & 6.18         & 0.04        & 1.66       & 2.&1.   \\\\  339 $^2$ &0.602& 5.07       &  -           & 1.38      &   -   &   -    \\\\ m11      &     &5.42         &  -           & 1.30        & 1.8 & 10.  \\\\ m12      &     &5.51        &  -           & 1.31       & 1.8 & 6. \\\\ m13     &      &5.54        &  -           & 1.33       & 1.8 & 2. \\\\ m14     &      & 5.55       &  -           & 1.34       & 1.8 & 1. \\\\ \\hline \\end{tabular}}  $^1$ from the Kakazu et al. sample ; $^2$ from the Xia et al. sample  \\end{table}  \\begin{table} \\centering \\caption{Comparison of model calculated O/H with Kakazu et al results} \\begin{tabular}{ccccccc} \\hline  \\hline \\  ID  &   z  & mod   & KCH1$^1$ & KCH2 $^1$    \\\\ \\hline \\ NB816&      &       &      &        \\\\ \\ 40   &0.629 &   8.82 &7.86 $<$ 8.03$<$8.25 & 7.84$\\pm$ 0.02 \\\\ \\ 76   &0.6319&  8.82  & $>$8.55& 7.89$\\pm$0.01  \\\\ \\ 118  &0.6439&  8.78  & 6.93$<$7.16$<$7.44   & 7.67$\\pm$0.03\\\\ \\ 195  &0.628 & 8.48   & 6.78$<$7.06$<$7.44   & 7.53$\\pm$0.03\\\\ \\ 252  &0.64  & 8.81   & 7.68$<$7.87$<$8.14   & 7.72$\\pm$0.01 \\\\ \\ NB912&      &         \\\\ \\ 3    &0.82   & 8.78  &7.26$<$ 7.43$<$7.65   & 7.71 $\\pm$0.02 \\\\ \\ 6    &0.83  &  8.62  & 7.40$<$7.68$<$8.14   & 7.79$\\pm$0.03 \\\\ \\ 9    &0.833 &  8.30  & $>$7.7 & 7.72$\\pm$0.02 \\\\ \\ 10   &0.829 &  8.82  &7.59$<$ 7.72$<$7.89   & 7.74$\\pm$0.02 \\\\ \\ 20   &0.820 &  8.78  &7.18$<$ 7.36$<$7.58   & 7.55$\\pm$0.02  \\\\ \\ 239  &0.8273&  8.81  & $>$7.34&7.43$\\pm$0.02 \\\\ \\ 270  &0.8176&  8.81  & 7.28$<$7.43$<$7.61   & 7.50$\\pm$0.02 \\\\ \\ 60 $^2$ &0.393  & 8.82 &7.67$<$ 7.96$<$8.57 & 7.81$\\pm$0.02 \\\\ \\hline \\end{tabular}  $^1$ Kakazu et al (2007, tables 4 and 5). KCH1:data obtained by direct method, KCH2:data based on Yin et al (2007) method.   \\caption{Comparison of model calculated O/H with Xia et al results} \\centering \\begin{tabular}{ccccccc} \\hline  \\hline \\ ID   &   z     & mod&  Xia et al $^1$  &   \\\\ \\hline \\ 339  & 0.602&  8.25&  8.10$^{+0.20}_{-0.16}$ &   \\\\ \\ 364  & 0.642&  8.30&  8.22$^{+0.16}_{-0.13}$  &  \\\\ \\ 246  & 0.696&  8.82&  7.71$^{+0.28}_{-0.27}$   &    \\\\ \\ 454  & 0.847&  8.82&  8.25$^{+0.23}_{-0.23}$   &  \\\\ \\ 258  & 0.998&  8.75&   7.49$^{+}_{-0.17}$  & \\\\ \\ 432  & 1.573&  8.79&  8.25$^{+}_{-0.26}$   &  \\\\ \\ 563  & 1.673&  8.79&  8.37$^{+}_{-0.28}$   &  \\\\ \\ 103  & 1.682&  8.79&   7.97$^{+}_{-0.22}$  &  \\\\ \\ 195  & 1.745&  8.7 &  8.38$^{+}_{-0.27}$   & \\\\ \\ 242  & 2.070&  8.66&  8.32$^{+}_{-0.29}$   &   \\\\ \\ 578  & 2.315&  8.66&  8.27$^{+}_{-0.26}$   &    \\\\ \\hline  \\end{tabular}  $^1$ Xia et al (2012, table 2).  Results obtained by R23  empirical calibrators.  \\caption{Comparison of model calculated O/H with Henry  et al results} \\centering \\begin{tabular}{cccccccc} \\hline  \\hline \\ ID   &   z      & mod$_{SB}$$^1$& mod$_{AGN}$$^2$ & KK04$^3$ & Maiolino$^4$ & FMR$^5$  \\\\ \\hline \\ 1    & 1.82     &8.82  & 8.6   & 8.16$^{+0.16}_{-0.15}$ & 7.59$^{+0.44}_{-0.16}$  & -     \\\\ \\ 2    & 1.73     &8.82  & 8.78  & 8.43$^{+0.16}_{-0.11}$ & 8.03$^{+0.33}_{-0.33}$  & 8.17    \\\\ \\ 3    & 1.77     &8.82  & 8.82  & 8.65$^{+0.14}_{-0.11}$ & 8.47$^{+0.14}_{-0.26}$  & 8.33   \\\\ \\ 4    & 1.74     &8.82  & 8.78  & 8.82$^{+0.08}_{-0.10}$ & 8.68$^{+0.09}_{-0.12}$  & 8.60    \\\\ \\hline  \\end{tabular}  \\flushleft  $^1$ calculated by models referring to the SB;  $^2$ calculated by models referring to the AGN  $^3$ calculated by the  Kobulnicky \\& Kewley (2004) calibration;  $^4$ calculated by the Maiolino et al. (2008)  calibration;  $^5$ the metallicity predicted from the local FMR (Mannucci et al 2012)  \\end{table}   \\begin{figure} \\includegraphics[width=9.0cm]{fig6.eps} \\caption{N/H versus O/H  from the sample of galaxies presented in Fig. 5. }  \\includegraphics[width=9.0cm]{fig7.eps} \\caption{O/H calculated by the direct method   and other empirical methods versus O/H calculated by detailed modelling. Kakazu et al  data : black circles (direct method); black +  (Yin et al method);  Xia et al data : blue crosses ; green symbols refer to the Henry et al data: large symbols result from AGN models The solid line represent the 1:1 line. } \\includegraphics[width=9.0cm]{fig8.eps} \\caption{The O/H relative abundances as a function of z. Kakazu et al   : black circles (direct method); black +  (Yin et al method); red squares (models). Xia et al data ; blue crosses; red diamonds (models);  Henry et al data (green symbols), red stars (models) } \\end{figure}   \\subsection{Comparison of the O/H results} Finally, we   compare in Tables 7, 8 and 9 the results for O/H calculated by the present  models and those determined by Kakazu et al., adopting the direct method and   by  Xia et al. and Henry et al. adopting  empirical calibrators. Our results  lead to O/H = 10$^{-4}$ -- 6.6 10$^{-4}$.  Actually, the physical parameters adopted by Kakazu et al and Xia et al. to characterize the emitting gas are different from those determined by the best fit of the line ratios in the present paper, e.g. Kakazu et al  fixed the electron density  to \\ne=100 \\cm3,  because the line ratios, which depend  directly on the density, were not observed. In our models \\ne ranges between $\\sim$ 600 and $>$ 1000 \\cm3   due to compression  (Fig. 2) and recombination downstream.  Moreover, Figs. 4  shows that  SB star  temperatures calculated for the Henry et al sample   are higher than those calculated from the Xia et al sample  and  higher than those calculated for a large sample of galaxies in Fig. 5. The ionization parameters calculated for  these stacked spectra are lower than those calculated for the Xia et al sample, but they are consistently located  throughout the large sample presented in Fig. 5.  Figs. 7 and 8 show that a  large gap appears   between the O/H results obtained by the direct method (and that of Yin et al 2007) and those obtained by detailed calculations, however, the trend of O/H with z is roughly similar for our results and those of Xia et al for  galaxies at z$\\geq$ 1.  The large gap between our O/H results  and those  obtained by the direct method  or by empirical calibrators has been explained in Sect. 2. Namely, the line intensities are calculated throughout a cloud integrating on  regions showing different \\Te,  \\ne and, consequently, different fractional abundances of the ions. Regions corresponding to relatively low \\Te  contribute mostly to low ionization  level lines, while a high \\Te contributes to relatively high ionization level lines. So the final [OIII] and [OII] line calculated intensities will be lower than those calculated  adopting the  optimum \\Te and \\ne. To reproduce the observed [OIII]/\\Hb line ratios  a higher O/H is then needed by the model.  Concluding, the relative O/H abundances  calculated by the direct method by Kakazu et al (2007) and empirical methods by Xia et al. (2012) and Henry et al. (2013) are  lower limits because they  adopted  physical conditions in the emitting nebulae  different from those consistently calculated by the detailed modelling.   The relatively high  O/H ratios calculated in this paper  reduce the low-metallicity  character of galaxies at higher z. Moreover, Figs. 4, 5 and 8   confirm  that  the critical  redshift for the scattering of   metallicity  started  at z $\\leq$ 1.  Finally,   in this paper which deals with the modelling of relatively high redshift galaxies  on the basis of  [OIII]/\\Hb and [OII]/\\Hb observations,  we claim that the  [OIII]5007+/\\Hb and [OII]3727+/\\Hb line ratios alone are not sufficient to constrain the results."
    },
    "1402/1402.4852_arXiv.txt": {
        "abstract": "We compute a large grid of photoionization models that covers a wide range of physical parameters and is representative of most of the observed PNe. Using this grid, we derive new formulae for the ionization correction factors (ICFs) of He, O, N, Ne, S, Ar, Cl, and C. Analytical expressions to estimate the uncertainties arising from our ICFs are also provided. This should be useful since these uncertainties are usually not considered when estimating the error bars in element abundances.  Our ICFs are valid over a variety of assumptions such as the input metallicities, the spectral energy distribution of the ionizing source, the gas distribution, or the presence of dust grains. Besides, the ICFs are adequate both for large aperture observations and for pencil-beam observations in the central zones of the nebulae.  We test our ICFs on a large sample of observed PNe that extends as far as possible in ionization, central star temperature, and metallicity, by checking that the Ne/O, S/O, Ar/O, and Cl/O ratios show no trend with the degree of ionization.  Our ICFs lead to significant differences in the derived abundance ratios as compared with previous determinations, especially for N/O, Ne/O, and Ar/O. ",
        "introduction": "\\label{sec:intro}  Planetary nebulae (PNe) are powerful tools to study chemical evolution in galaxies. The spectra of PNe have bright emission lines that allow us to calculate physical conditions and ionic abundances in the ionized gas. The abundance of sulfur, neon, argon, and chlorine in PNe reflect the abundances of the interstellar medium at the time the progenitor stars were born. Therefore, these abundances can be used as tracers of interstellar abundances. On the other hand, helium, nitrogen, carbon, and oxygen abundances change as a consequence of nucleosynthesis processes occurring in the star before and during the PN phase. The analysis of these abundances can give us clues on the efficiencies of nucleosynthesis in low and intermediate mass stars.  The total abundance of one particular element is obtained by adding up the ionic abundances of all the ions present in a nebula. However, not all the ions are observed, either because the lines are emitted in a different spectral range from the one observed or because the lines are too weak to be detected. Thus, one must estimate the contribution of unobserved ions to the total abundance using ionization correction factors (ICFs). First ICFs were proposed by \\citet{Peimbert_69} and they were based on ionization potential similarities. However these ICFs should be taken with caution since the ionization structures do not depend only on ionization potentials \\citep[see, e.g.,][]{Stasinska_02}. Alternatively, ICFs may be derived from photoionization models, as those derived, e.g., by \\citet{Stasinska_78}; Kingsburgh \\& Barlow (1994, hereafter KB94); and \\citet{Kwitter_01}. They are more reliable but so far have been based on a small number of models.  In this paper we derive new recipes for the ICFs of He, O, N, Ne, S, Ar, Cl, and C using a large grid of models that spans over a wide range of parameters and is representative of most of the observed PNe. We also provide analytical formulae to estimate the error bars associated with our ICFs. Our motivations are two. First, some of the commonly adopted ICFs do not correct properly for unobserved ions. Second, the errors in ICFs are usually not considered when estimating errors in element abundances. Our results should be useful for empirical abundance studies. In particular, an accurate determination of the element abundances and their associated uncertainties is essential in the study of abundance gradients in galaxies.  We examine the accuracy of our ICFs by applying them to several families of photoionization models computed from different assumptions. The studied performed by \\citet{Alexander_97} is the only one in the literature investigating the errors arising from the use of ICFs in PNe. Our sample of models cover a wider range of physical parameters and we include several assumptions not considered in \\citet{Alexander_97}, such as the non-constant density distribution, the combination of two models, or the inclusion of dust grains in the nebulae. We also test our ICFs on a large sample of observed PNe, and compare the computed abundances with the ones obtained with other ICFs (mainly the ones by KB94).  The analysis here is restricted to ions with emission lines in the optical range, ICFs for ultraviolet and infrared observations will be discussed in a forthcoming paper. ",
        "conclusions": "\\label{sec:concl}  We have computed a large grid of photoionization models covering a wide range of physical parameters representative of most of the observed PNe. Using this grid of models, we derived new recipes for the ICFs of He, O, N, Ne, S, Ar, Cl, and C that should be used for optical observations. We also provide analytical expressions for the associated uncertainties. This should be useful for empirical abundance studies since errors in ICFs are usually not considered when estimating errors in element abundances.  The ICFs are valid for all the computed families of models, which include models with different input metallicities, different gas density distributions, presence or absence of dust, or different spectral energy distribution of the ionizing source. Besides, our ICFs are valid for volume integrated observations as well as for pencil beam observations in the vicinity of the central star. Therefore they can be safely used in a large variety of PNe observations.  For each correction scheme, we define the region of validity, outside this zone the error bars are so large that no meaningful abundance can be computed. Objects presenting such \\gihe\\ and \\gio\\ should be removed from abundance computations using our ICFs for the corresponding element. A tailored photoionization model is recommended for such cases.  In the case of helium we do not suggest the derivation of an ICF based on other ions since the relative populations of helium ions depend essentially on the \\efftemp\\ of the central star while the ones of metal ions also depend on the ionization parameter. We recommend to use ICF(He$^{+}$ + He$^{++}$) = 1 and, for those objects with $\\omega \\lesssim 0.45$, incorporate the correction for unobserved ions in the error bars.  For nitrogen and neon, we propose ICFs based both on \\gio\\ and \\gihe. The uncertainties associated with our computed N/O and Ne/O values are significantly lower (especially for PNe with low \\efftemp\\ central stars) than those derived using the ICFs suggested by KB94.  As for the other elements we suggest ICFs based on \\gio\\ or \\gihe. In general, the uncertainties related to our ICFs are significantly lower than the ones related to other ICFs in the literature.  We have shown that the error bars associated with ICFs should be taken into account in the total error budget on abundances since they represent a significant budget of total errors abundances.  Although our sulfur abundances are somewhat higher than those derived from the ICF by KB94, the S/O values obtained in most of the studied PNe are still lower than those in Orion or in the Sun. This suggests that the ICF is not the principal cause of the sulfur anomaly.  The abundance ratios Ne/O, S/O, Cl/O, and Ar/O computed with our ICFs in a wide sample of observed PNe (covering a wide range of ionization, central star temperature, and metallicity) do not show any noticeable trend with the degree of ionization. This validates our ICFs for abundance studies.  The ICFs defined in this paper will be included in the next release of the Pyneb package."
    },
    "1402/1402.1646_arXiv.txt": {
        "abstract": "{ Recently discovered scattered light at 3-5 $\\mu$m from low-mass cores (so-called \"coreshine\") reveals the presence of grains around 1 $\\mu$m, which is larger than the grains found in the low-density interstellar medium. But only about half of the 100+ cores investigated so far show the effect. This prompts further studies on the origin of this detection rate. } { We aim to supply criteria for detecting scattered light at 3.6 $\\mu$m from molecular cloud cores. } { From the 3D continuum radiative transfer equation, we derive the expected scattered light intensity from a core placed in an arbitrary direction seen from Earth. We use the approximation of single scattering, consider extinction up to 2nd-order Taylor approximation, and neglect spatial gradients in the dust size distribution. We analyze how scattered light can outshine the absorbing effect of extinction in front of background radiation by the core for given grain properties, anisotropic interstellar radiation field and background field. The impact of the directional characteristics of the scattering on the detection of scattered light from cores is calculated for a given grain size distribution, and local effects like additional radiation field components are discussed. The surface brightness profiles of a core with a 1D density profile are calculated for various Galactic locations, and the results are compared to the approximate detection limits. } { We find that for optically thin radiation and a constant size distribution, a simple limit for detecting scattered light from a low-mass core can be derived that holds for grains with sizes smaller than 0.5 $\\mu$m. The extinction by the core prohibits detection in bright parts of the Galactic plane, especially near the Galactic center. For scattered light received from low-mass cores with grain sizes beyond 0.5 $\\mu$m, the directional characteristics of the scattering favors the detection of scattered light above and below the Galactic center, and to some extent near the Galactic anti-center. We identify the local incident radiation field as the major unknown causing deviations from this simple scheme. } { The detection of coreshine at a wavelength of 3.6 $\\mu$m is a complex interplay of the incident radiation, the extinction of the background radiation, the grain properties, and the core properties like sky position and mass. } ",
        "introduction": "Dense regions in molecular clouds have been identified as the site where star formation starts. As the initial conditions impact the formation process, much effort has been devoted to analyze the change in conditions from the diffuse interstellar medium (ISM) to these dense cold cores \\citep[for an overview, see, e.g., ][]{2007ARA&A..45..339B}. Through absorption, scattering and emission, cosmic dust grains embedded in the core gas enshroud the view on the star formation process at optical wavelengths, but also provide information on the physical conditions at longer wavelengths.  A key quantity of the grains is their size distribution. For the diffuse ISM, \\citet{1977ApJ...217..425M} were the first to suggest a power-law distribution with a size limit of 0.25 $\\mu$m for non-graphite grains (hereinafter called \"MRN(0.25)\"). They emphasized that only weak constraints could be derived for the larger grains in the distribution. For the denser phases of the ISM, however, the modeling of thermal emission and extinction of grains provided indirect evidence for the existence of a population of grains beyond MRN(0.25) \\citep[e.g.][]{ 2003A&A...398..551S, 2006MNRAS.373.1213K, 2006A&A...451..961R, 2008ApJ...684.1228S, 2009ApJ...690..496C, 2009ApJ...699.1866C, 2009ApJ...693L..81M}. \\citet{1996A&A...309..570L} modeled the J, H, and K band scattered light images of the Thumbnail nebula and concluded that the limiting grain size exceeds 0.25 $\\mu$m.  In a recent paper, evidence in the same direction was reported for the core LDN183 by \\citet{2010A&A...511A...9S} based on the detection of a weak extended surface brightness excess with respect to the large-scale infrared background from the core visible in the 3.6 and 4.5 $\\mu$m IRAC bands of the {\\it Spitzer} space telescope called \"coreshine\". Since the core is too cold for large grain thermal emission, and the emission of stochastically heated particles like polycyclic aromatic hydrocarbons or very small grains should be restricted to the outer parts of the core to receive enough UV/optical energy, the radiation was interpreted as scattered interstellar light. Furthermore, despite the absence of strong stochastically heated emission in the 4.5 $\\mu$m band, surface brightness excess was seen also in this band. The interpretation as scattered light required a scattering efficiency of the dust that is only provided by grains with radii up to 1 $\\mu$m.  Although the central core of LDN183 is surrounded by slightly more mass than most low-mass cores, its general conditions were not substantially different from that of many other low-mass cores. Therefore, it was not entirely surprising that an investigation of other cores in the IRAC bands led to the finding that about half of the selected cores showed scattered light \\citep{2010Sci...329.1622P} \\citep[see also][for a small subsample of cores with 3.6 $\\mu$m excess]{2009ApJ...707..137S}.  A radiative transfer analysis of model cores revealed that the scattered light morphology (the morphology referring to the gradients across the core, not the total surface brightness) of low-mass cores should show no strong selection effect with position in the Galaxy aside from a weak enhancement towards the Galactic center (GC). For more massive cores with 10 $M_\\odot$, the morphology showed a crescent shaped asymmetry in the morphology of the surface brightness excess. A variation of the total expected flux as a function of sky position was not investigated.  For a special region, namely the Gum/Vela region, \\citet{2012A&A...541A.154P} found a drastically reduced occurrence of scattered light (3 detections plus 3 uncertain cases out of 24 objects). They did not attribute this absence to systematic detection difficulties but to the action of a supernova remnant blast wave which has affected the grain size population in this region. The analysis of scattered light based on IRAC photometry is compromised by PAH emission that in principle can contaminate all four IRAC bands. This is true for any star forming region like $\\rho$~Oph or Vela that is irradiated by close medium or high-mass stars. When scattered light and PAH emission regions are spatially separated, one can use the almost PAH-free 4.5 $\\mu$m band to distinguish or use that the PAHs appear in all four IRAC channels opposite to the scattered light being strongest in the 3.6 $\\mu$m band and decreasing in flux with wavelengths. When a region shows both scattered light and PAH emission, it is difficult to disentangle them.  The work presented in this paper focuses on deriving the general conditions for the detection of scattered light at wavelength of 3.6 $\\mu$m from a low-mass molecular cloud core. Assuming various grain size limits, we explore the different factors influencing the detection. In Sect.~2 we derive approximative formula for the coreshine detection limit from the radiative transfer equation. In Sect.~3 we compare the formula with results from a grid of 3D radiative transfer calculations for a model core spanning an all-sky map of coreshine surface brightness profiles. Observational constraints are discussed in Sect.~4 for a set of example cores, and we conclude our findings in Sect.~5. ",
        "conclusions": "\\label{conclusions} We find that the detection of scattered light at 3.6 $\\mu$m from low-mass molecular cloud cores is both influenced by local effects like the radiation field in the vicinity of the core and global effects like the interplay of core position with respect to the GC and the directional characteristics of the scattering. For optically thin radiation and a constant grain size distribution throughout the core, we derive a simple limit for detecting coreshine that holds for grains with sizes smaller than 0.5 $\\mu$m. The extinction of background radiation by the core prohibits detection in large parts of the Galactic plane and especially near the LoS towards the GC. For grain distributions extending to 1 $\\mu$m, the forwardly-peaked directional characteristics of the scattering favors the detection of coreshine near the GC. After combining the two effects, we identify regions above and below the GC and anti-center to be those with enhanced detection. Two further quantities enter the scattering-extinction balance formula: the local background and the local ISRF. In this work we use the only all-sky map having an absolute flux calibration which is the DIRBE map. Due to the coarse Healpix grid it provides only a mean value where strong deviations can occur locally. Moreover, it contains the stellar contribution which is not included in the core background modeling. Another uncertainty is the local incident radiation field that is observed in star formation regions and that is likely produced by the warm dust near the forming stars. Its directional variation, flux, and spectrum is unknown however, and very difficult to measure. We expect that only individual core modeling may derive constraints on the local field: for the core LDN260 modeled in \\citep{2013A&A...559A..60A} an enhancement of the incident radiation field was necessary to reproduce the observed coreshine. The derived scheme can be used to identify cores which deviate from the expected flux and might be of interest in terms of nearby radiation sources, PAH contamination, or grain properties. Combining the global constraints on coreshine detection, cores above or below the GC are best detectable. The moderate amplification by backward scattering of GC radiation along with the lower background supports detection of coreshine near the anti-center.  The opacities by \\citet{1984ApJ...285...89D} used in this paper have served as reference data for decades as much as later work showed in how much modification in the chemical composition, shape, size distribution will alter the values. It remains true however that we have little observational information on the exact properties of larger grains which leaves ample room for changing the dust properties and distribution and arriving at the same effects. Therefore, instead of investigating the impact of various dust model assumptions on the derived global effects, we expect progress from modeling multi-wavelength data for a large number of single cores as presented in \\citet{2010A&A...511A...9S} and \\citet{2013A&A...559A..60A} where dust modeling is better constrained. Future work also will include results from coagulation calculations \\citep{L1506C}."
    },
    "1402/1402.6911_arXiv.txt": {
        "abstract": "The recent measurement of two solar mass pulsars has initiated an intense discussion on its impact on our understanding of the high-density matter in the cores of neutron stars. A task force meeting was held from October 7-10, 2013 at the Frankfurt Institute for Advanced Studies to address the presence of quark matter in these massive stars. During this meeting, the recent observational astrophysical data and heavy-ion data was reviewed. The possibility of pure quark stars, hybrid stars and the nature of the QCD phase transition were discussed and their observational signals delineated. ",
        "introduction": "The recent observations of the pulsars PSR J1614-2230 and PSR J0348+0432 proofs beyond doubt the existence of neutron stars with masses close to two solar masses \\cite{Demorest:2010bx,Antoniadis:2013pzd}. Moreover, there are hints that there might exist even more massive neutron stars, e.g. from a recent analysis of the black-widow pulsar PSR J1311-3430 \\cite{Romani:2012rh} and pulsar PSR J1816+4510 \\cite{Kaplan:2013hii}.  In addition, future astrophysical instruments will set much more refined constraints on the properties of neutron stars. The future Five hundred meter Aperture Spherical Telescope FAST and the Square Kilometre Array SKA will be sensitive enough to find an order of magnitude more millisecond pulsars and pulsars in binary systems than those found up to now, increasing the number of pulsar mass measurements accordingly \\cite{Smits:2009vd,Cordes:2005gp}. The gravitational wave detectors advanced LIGO and advanced VIRGO, both starting operation in 2014, will have an order of magnitude increase in their sensitivity to find gravitational waves from compact binary mergers with a projected rate of measuring about several events per year \\cite{Abadie:2010cf}. Such increased sensitivity will allow scientists to determine the basic properties of binary neutron star systems to within a few percent \\cite{Aasi:2013jjl}. Finally, future X-ray satellite missions such as NICER \\cite{NICER} and LOFT \\cite{Feroci:2011jc,Bozzo:2011vn,Bozzo:2013txa} plan to measure the radius of compact stars to unprecedented precision.  In reaction to these astrophysical observations and future prospects, it seems timely to discuss the role of quark matter in the cores of neutron stars. We do that by investigating quark matter properties that are relevant for compact stars and the possible impact of quark matter on astrophysical observations. A rapid reaction task force meeting of the ExtreMe Matter Institute (EMMI) on ``Quark Matter in Compact Stars'' was held from October 7-10, 2013 at the Frankfurt Institute for Advanced Studies (FIAS) in Frankfurt/Main, Germany. During the task force meeting, the properties of dense quark matter relevant for compact star physics were discussed in intense closed working sessions. A previous meeting of this kind focussing on the role of hyperons in neutron stars and the problem of low maximum masses of compact stars in microscopic models with hyperons, Hyperons2012, was held in Warsaw in November of 2012.  Some of the key questions addressed in the task force meeting on quark matter in compact stars were \\begin{itemize}  \\item what are the properties of dense matter in the core of massive pulsars with two solar masses?  \\item can one constrain effective models of dense QCD matter?  \\item is there still room for quark matter to be present in neutron stars?  \\item would the observation of a say 2.4 solar mass neutron star exclude exotic matter in its core?  \\item is there a smoking gun for the presence of quark matter in compact stars?  \\item are quark stars still a viable possibility? Can quark stars and hadronic stars coexist?  \\item what can we learn from radius measurements about the properties of QCD matter at extreme densities?  \\end{itemize} The theoretical investigation of these questions is related to two main research themes: the structure and dynamics of neutron star matter, and the properties of the quark-gluon plasma and the phase structure of strongly interacting matter. The properties of dense matter in the core of a neutron star determine the global properties of compact stars, i.e.\\ the mass-radius relation, which was a key issue of the task force meeting. This is also relevant for the investigation of the QCD phase diagram and its possible phase transitions at high densities and low temperatures. Possible cross relations of these two research areas were unavoidably part of the scientific discussion of the task force meeting. The presentations are available at the task force website at https://indico.gsi.de/conferenceTimeTable.py?confId=2420\\#all.detailed. ",
        "conclusions": ""
    },
    "1402/1402.4313_arXiv.txt": {
        "abstract": "We propose a new method to analyse the alignment of optical polarisation vectors from quasars. This method leads to a definition of intrinsic preferred axes and to a determination of the probability $p^{\\sigma}$ that the distribution of polarisation directions is random. This probability is found to be as low as $3.0\\, 10^{-5}$ for one of the regions of redshift. ",
        "introduction": "\\label{sec:Intro} Fifteen years ago, Hutsem\\'ekers discovered that the optical polarisation vectors from distant quasars are not randomly distributed \\cite{key-1}. The original study considered a sample of 170 quasars, and the probability of a random distribution was 0.5 \\%. Since then, further measurements have been added \\cite{key-2,key-10,key-3}, and the significance of the effect has grown: with the present sample of 355 polarisation data points, the alignment effect has been estimated to have a probability to be random lower than 0.1 \\% in some regions of the sky defined by cuts on right ascension and redshift.  The original test however had a weak point, as the angles of the polarisation vectors are measured with respect to their local meridian. Therefore, the strength of the alignment, and the significance of the effect, depend on the choice of the spherical coordinate system, so that it is possible to find an axis for which the effect becomes negligible, i.e. for which the data look random. Indeed, the original statistical tests do not take into account the fact that polarisation vectors come from different lines of sight and thus are not defined in the same plane. To remedy this problem, \\cite{key-6} have proposed to parallel-transport the polarisation vectors along geodesics of the celestial sphere onto one point in order to compare them. While the distribution of polarisation angles still depends on the point where it is built, it does not depend on the coordinates, and one can build robust statistical tests (see Jain et al. (2004) and Hutsem\\'ekers et al. (2005)). Unfortunately, these tests have to resort to the generation of a large number of random data in order to determine the significance of the signal and do not lead to a clear characterization of the effect. Furthermore, to make an intrinsic test, polarisation vectors must be transported to the location of the quasars. However, this is an arbitrary choice as they can be transported to any point on the celestial sphere. The sum of the parallel-transported polarisation vectors then makes a continuous axial vector field on the sphere, and by the hairy ball theorem \\cite{hbt}, it is still possible to find at least two points where the effect vanishes. Hence the quantification of the size of the effect seems uncertain again.  We propose here another method which is totally independent of the coordinate system, and which quantifies unambiguously the alignment effect. It allows us to compare the polarisation vectors of sources located at different angular coordinates and leads to the characterization of the effect through a blind analysis of the data. The basic idea is to consider the physical polarisation vectors as 3-dimensional objects rather than 2-dimensional ones embedded in their polarisation plane. These 3-dimensional objects are the directions of the electric field oscillations and they are the physical objects which are measured. As we are dealing with a number of vectors, it is clear that it will be possible to define preferred directions. We can also use the same method to study the dependence on redshift, position in the sky or degree of linear polarisation by imposing cuts on these variables and repeating the study for the corresponding sub-sample.  We devote the present paper to the construction of this new statistical method applied here to the analysis of quasar polarisation data, in which the evaluation of the significance of the signal will be largely analytic. The second section of this paper explains the details of our statistical method. It also contains illustrations and discussions related to the statistical background, when polarisation angles are assumed to be uniformly distributed, and is compared with the original study of \\cite{key-3} using the same cuts. The third section is devoted to the application of this new statistical method first globally, then to slices in redshift. We also consider there the dependence of the alignment on the various parameters, the possibility of a cosmological alignment, and define our final most significant regions exhibiting an anomalous alignment of polarisation vectors. ",
        "conclusions": "\\label{sec:Ccl} We have presented in this paper a new coordinate-invariant method to detect polarisation alignments in sparse data, and applied it to the case of alignments of optical polarisation vectors from quasars. We showed that we automatically recover regions previously found, and we refined their limits based on objective criteria (see Table~\\ref{tab:best}). As a byproduct, the directions of alignments in space are unambiguously determined. The method we propose is powerful, as the coordinate-invariant significance levels are semi-analytically determined. The remaining drawback is that the determination of the global significance levels relies on very time-consuming Monte-Carlo simulations. We believe that this new analysis puts the alignment effect on stronger grounds as the global significance level is as low as $3.0\\, 10^{-5}$ for some regions of redshift.  However, one has to note that the significance levels obtained in this papers and those reported in \\cite{key-3} are not in full agreement. Indeed, the Z-type tests used in \\cite{key-1} and \\cite{key-6} reshuffle the measured polarisation directions while keeping the source locations to evaluate the background. The advantage is that any systematic effect vanishes automatically through this method. The disadvantage is that it washes out global effects, or alignments present for a large number of quasars. Random generation of polarisation angles, as used here or in \\cite{key-6}, has the opposite features: we can detect global alignments, but we are sensitive to systematic effects. Hence the two methods do not need to be in full agreement. One should note, however, that the sample of optical polarisation measurements of quasars considered here comes from many independent observational campaigns, so that a common bias is unlikely (see \\cite{key-1} and \\cite{key-3} for discussion). Furthermore, \\cite{key-6} have addressed this question of global systematic effect with the sample of 213 quasar polarisation measurements available at the time by comparing analyses of Z-type tests using uniform polarisation angle distribution and distribution made by reshuffling. They found that the significance level could decrease by factors of the order three. We expect this to be the case in this study. Note that this uncertainty is of the same order as that we estimated using the Jackknife algorithm.  Applying our method to the sample of 355 quasars compiled in \\cite{key-3}, we identified the following main features. The directions of alignments show a dependence on the redshift of the sources. Although this dependence seems discontinuous, one should note that we detected significant alignments for redshift intervals where the distribution of data peaks. Thus, more data in regions of redshift with poor statistics are required in order to study this dependence in more details. As seen in Fig.~\\ref{fig:finestruct}, for a given redshift interval, alignment seems to be mainly due to quasars well localized toward specific directions of the sky. Furthermore, no strong evidence has been found for a dependence on the degree of linear polarisation.  As a result of the application of our new method to the present sample of optical polarisation measurements of quasars, and in agreement with \\cite{key-3}, we found several distinct sub-samples of sources well localized in redshift and position on the sky that show unexpected alignments of their polarisation vectors. We established two regions towards the North galactic pole, one at low and the other at high redshift, and only one towards the South galactic pole at intermediate redshift, which possibly dominates the whole southern sky. Besides the regions previously detected, or their improved version, we also showed that there exists the possibility of a cosmological alignment."
    },
    "1402/1402.2229_arXiv.txt": {
        "abstract": "We studied the total magnetic field strength in normal star-forming galaxies estimated using energy equipartition assumption. Using the well known radio--far infrared correlation we demonstrate that the equipartition assumption is valid in galaxies at sub-kpc scales.  We find that the magnetic field strength is strongly correlated with the surface star formation rate in the galaxies NGC 6946 and NGC 5236.  Further, we compare the magnetic field energy density to the total (thermal + turbulent) energy densities of gas (neutral + ionized) to identify regions of efficient field amplification in the galaxy NGC 6946.  We find that in regions of efficient star formation, the magnetic field energy density is comparable to that of the total energy density of various interstellar medium components and systematically dominates in regions of low star formation efficiency. ",
        "introduction": "\\label{s:intro} \\vspace{-5mm}  It is believed that the magnetic field in star forming galaxies are amplified and maintained by dynamo action. At small scales ($\\lesssim1$ kpc), the turbulent dynamo action, driven by Supernova shocks and star formation, amplifies the magnetic field strength efficiently up to energy equipartition values in $\\sim 10^6 - 10^7$ years.  Such regions are characterized by small-scale random field with low degree of polarization \\citep[see e.g.][]{brand12, arsha09, bhat13}.  At larger scales ($\\sim1-10$ kpc), the mean field dynamo action is responsible for amplification of the field in $\\sim10^8-10^9$ years giving rise to large-scale coherent field with high degree of polarization \\citep{arsha09, subra08, shuku06}.  Magnetic field strength grows due to stretching and twisting of the field lines up to energy equipartition values, i.e., energy density in magnetic field is similar to the kinetic energy density due to turbulent motions of gas.  Numerical magnetohydrodynamic simulations of turbulent interstellar medium (ISM) have revealed that the magnetic field ($B$) and gas density ($\\rho_{\\rm gas}$) are coupled as $B\\propto\\rho_{\\rm gas}^\\kappa$, where $\\kappa$ is the coupling index \\citep{cho00, fiedl93, grove03}. Under the condition of equipartition of energy, $\\kappa$ assumes the value 0.5 \\citep{cho00}. It is therefore expected that the radial scale length of the magnetic field energy and that of gas density should remain similar.  However, in contradiction, it is often observed that the magnetic energy density dominates over the total energy density of gas especially towards the outer parts of the galaxies \\citep{basu13, beck07}. Star formation activity and shocks arising from supernova explosions drives turbulence at the smallest scale in the ISM.  It is therefore imperative to do spatially resolved study of the various competing forces in the ISM with the magnetic field strength based on star formation activity. ",
        "conclusions": ""
    },
    "1402/1402.2645_arXiv.txt": {
        "abstract": "We present results from 3D magnetohydrodynamic (MHD) simulations of the emergence of a twisted convection zone flux tube into a pre-existing coronal dipole field. As in previous simulations, following the partial emergence of the sub-surface flux into the corona, a combination of vortical motions and internal magnetic reconnection forms a coronal flux rope. Then, in the simulations presented here,  external reconnection between the emerging field and the pre-existing dipole coronal field allows further expansion of the coronal flux rope into the corona. After sufficient expansion, internal reconnection occurs beneath the coronal flux rope axis, and the flux rope erupts up to the top boundary of the simulation domain ($\\sim$ 36 Mm above the surface). We find that the presence of a pre-existing field, orientated in a direction to facilitate reconnection with the emerging field, is vital to the fast rise of the coronal flux rope. The simulations shown in this paper  are able to self-consistently create many of the surface and coronal signatures used by coronal mass ejection (CME) models. These signatures include: surface shearing and rotational motions; quadrupolar geometry above the surface; central sheared arcades reconnecting with oppositely orientated overlying dipole fields; the formation of coronal flux ropes underlying potential coronal field; and internal reconnection which resembles the classical flare reconnection scenario. This suggests that proposed mechanisms for the initiation of a CME, such as ``magnetic breakout'', are operating during the emergence of new active regions. ",
        "introduction": "Coronal mass ejections (CMEs) are among the most energetic phenomena associated with solar activity and space weather. These giant eruptions of solar  plasma and magnetic field are due to the sudden release of energy built up in the complexity of the magnetic fields in the solar atmosphere. Many CMEs are associated with {filament channels}, where the magnetic fields are strongly sheared, and hence are strongly non-potential and have significant free energy \\citep{2000JGR...10523153F, Klimchuk2001, LintonM2009}.  Twisted magnetic flux ropes are thought to play a major role in the onset and evolution of CMEs. In quiet Sun regions, pre-eruption prominences are interpreted as twisted coronal flux ropes \\cite[e.g.,][]{2010ApJ...724.1133G}. For active region CMEs, line of sight observations are more difficult and the scenario is not so clear, hence the use of idealized models is required to fully understand the role of flux ropes in these active region CMEs.  Almost all current CME models which include a magnetic field require either a pre-formed coronal flux rope  \\cite[e.g.,][]{2003ApJ...588L..45R,2005ApJ...630L..97T,2008ApJ...684.1448M}, or the formation of a flux rope from a highly sheared active region prior to or during the onset of eruption \\cite[e.g.,][]{1998ApJ...502L.181A,1999ApJ...510..485A,2000ApJ...529L..49A, 2003ApJ...585.1073A,2008ApJ...683.1192L}. Hence these models rely on either the transfer of sheared, non-potential field from beneath the surface, or the evolution of potential coronal magnetic field into non-potential field by shearing and/or rotational surface motions, as well as flux cancellation or magnetic diffusion. Recent simulations of flux emergence have shown that the partial emergence of a sub-surface twisted flux tube into the solar atmosphere leads to shearing motions \\citep{2004ApJ...610..588M} and sunspot vortical motions \\citep{2009ApJ...697.1529F}, and observations of sunspot rotations have also been interpreted as signatures of twisted flux tube emergence \\citep{2013SoPh..282..503K}. The \\textit{ad hoc} surface motions utilized by some CME models are motivated by observations of active regions, and the key features of these observations, such as shearing and rotation, are most likely a consequence of the emergence of twisted magnetic flux from beneath the surface. Hence it can be argued that these CME models rely on the emergence of twisted magnetic flux from the convection zone. However they do not self-consistently calculate a process for this flux emergence, as they do not include the lower solar atmosphere and convection zone, but instead use boundary conditions which have features that are associated with the emergence of new twisted flux.%  Early 3D simulations of flux emergence found that a twisted, buoyant, convection zone magnetic flux tube only partially emerges, with the axis confined to less than ten pressure scale heights (1.5 Mm) above the surface \\cite[e.g.][]{2001ApJ...554L.111F,2001ApJ...549..608M}. Later simulations found that a new flux rope structure forms in the corona, and the formation mechanism was attributed to either shearing and rotational motions observed at the surface \\cite[e.g.,][]{2009ApJ...697.1529F,Leake2013_1}, or magnetic reconnection \\cite[e.g.,][]{2004ApJ...610..588M,Archontis_Hood_2012}. For simulations without any coronal field, this rope expands and rises into the domain with speeds up to 33 km/s \\citep{2004ApJ...610..588M,2009ApJ...697.1529F}. For simulations with a pre-existing straight, constant-strength, field localized in the corona and aligned favorably for magnetic reconnection with the emerging field, the confining field is removed which allows a faster escape with speeds up to 60 km/s \\citep{2008A&A...492L..35A,2009A&A...507..995M,2010A&A...514A..56A,Archontis_Hood_2012}.  In this paper we address the scenario of how an eruptive flux rope can be formed in the corona. We study the emergence of twisted convection zone magnetic field into the corona and its interaction with a pre-existing dipole active region magnetic field, a field that is more complex than the spatially independent fields used in the previous studies mentioned above. We also address the issue of whether dynamical flux emergence of sheared field from the convection zone can capture the signatures and driving conditions used by CME models such as the so-called ``magnetic breakout'' model. This model relies on shear and /or rotational motions to create a sheared arcade from a potential field, and reconnection between different flux systems to initiate an eruption \\citep[e.g.,][]{1998ApJ...502L.181A,1999ApJ...510..485A, MacNeice04,2008ApJ...683.1192L}. In \\citet{2010ApJ...722..550L}, we performed 2.5D simulations of the emergence of twisted magnetic flux from the convection zone into various coronal configurations, such as simple dipole fields and the quadrupolar fields used in the magnetic breakout model. We found that in 2.5D the emergence process is unable to create an unstable configuration in the corona, due to the suppression of the emergence by dense plasma trapped on the emerging field. Further 2.5D simulations by \\citet{2013ApJ...764...54L} found that the slippage of magnetic field though the partially ionized regions of the solar atmosphere is a viable mechanism for allowing more magnetic flux to emerge, but still does not create unstable configurations. We therefore concluded that 3D motions are the only remaining plausible mechanism for this paradigm. It was shown by \\citet{2009ApJ...697.1529F}, and confirmed in our simulations of \\citet{Leake2013_1} (hereafter known as Paper I), that during 3D simulations of flux emergence, vortical motions of sunspots, driven by gradients in twist, are capable of twisting the field in the corona, creating a coronal flux rope. Therefore in this paper, we extend the magnetic breakout explorations of \\citet{2010ApJ...722..550L} to 3D, and attempt to drive the eruption of a coronal flux rope by emerging a  twisted flux tube into a pre-existing coronal dipole field in a 3D geometry. This pre-existing dipole is designed to represent the decaying field of an old active region. {\\color{blue}  Previously \\citet{Roussev_2012} have performed an emergence and eruption simulation in a similar configuration, on a global scale. Current computational resources make the simulation of dynamical emergence and eruption on a global scale difficult. Therefore in the study of \\citet{Roussev_2012} the surface signature from a simulation of flux emergence into a field-free corona, performed on the same scale as the simulations in this paper, 50 Mm, was used to drive the corona of a global simulation with a dipole field, on the scales of 500 Mm. In order to use the surface data from the small-scale flux emergence simulation to drive the coronal global simulation, three main assumptions were used. First, the length scale of the driving data was increased by an order of magnitude to match that of the global simulation. Second, the magnitudes of the non-force-free field, plasma pressure, and density, were reduced by 2, 6, and 9 orders of magnitude, respectively, to values representative of the corona, where the magnetic field should be nearly force-free. Third, it was assumed that the pre-existing coronal dipole field of the driven simulation did not need to be included in the driving simulation. This simulation produced an eruption-unstable coronal flux rope, which is our goal here, but the roles that these various non--self-consistent assumptions played in the dynamics is unknown. To explore this mechanism in detail, in a self-consistent set of simulations, we therefore restrict ourselves to a simulation length scale of 50 Mm. This allows us to simulate the entire domain in a single simulation.} ",
        "conclusions": "We have performed simulations of the emergence of convection zone flux tubes into the solar atmosphere with a pre-existing dipole coronal field, focusing on the interaction between emerging flux and dipole field, and the resulting dynamics. These % CME simulations  include both the dynamic emergence of magnetic flux into the corona and a dipole background coronal field representing a line-tied decaying active region. This is an improvement of the recent simulations of \\citet{Archontis_Hood_2012} which included a spatially independent, constant-strength, coronal field. In those simulations, the likelihood of eruption increased as the angle between the coronal field and emerging field became more favorable to external reconnection. From the results of our Paper I and this paper, we also conclude that the external reconnection is vital to the eruption process, which only happened when the dipole was aligned to maximize external reconnection.  Initially the azimuthal flux of the convection zone flux tube is much larger than that of the dipole and so the initial rise and emergence in the low atmosphere (photosphere/chromosphere) is similar to one where no dipole field exists. However, as the flux tube only partially emerges, not all its azimuthal flux is able to emerge into the corona. At a certain point in the emergence, the azimuthal fluxes of the two flux systems in the corona (emerging field and dipole field) are sufficiently balanced to make magnetic reconnection important.  The resulting atmospheric magnetic field (above $z=0$) closely resembles the quadrupolar geometry commonly used in the magnetic breakout model, with a central sheared arcade, side lobes and overlying dipole field. The fact that the simulations self-consistently produce this configuration without the need for kinematic boundary conditions is a step forward in improving the realism of the magnetic breakout model.  External reconnection between the emerging field and the dipole field above the central arcade allows further vertical expansion, while the horizontal expansion is suppressed by the formation of the side lobes due to the reconnection above. During the continued vertical expansion into the corona, equilibration of twist along emerged fieldlines causes apparent sunspot rotations. At present we cannot tell how much these sunspot rotation contribute to the accumulation of helicity and magnetic energy in the corona, relative to the other horizontal motions present during the flux emergence (e.g., shear flows and siphon flows), and relative to the vertical motions. This equilibration of twist effectively twists up the fieldlines in the corona, distorting the original sheared arcade structure.  % After a certain amount of vertical expansion, there is a noticeable build up of current density beneath the original flux tube axis and signs of internal (or flare-like) reconnection. Two O-points are formed above and below the  X-point of the reconnection site. The former is the intersection of a new coronal flux rope axis and the $y=0$ plane. We propose that both mechanisms (twisting and reconnection) contribute to the formation of a coronal flux rope.%  In these simulations, there is no eruption unless there is a pre-existing coronal dipole field, aligned to maximize external reconnection between emerging field and coronal field. Not only is the dipole field critical to the likelihood of eruption, but the ratio of dipole flux to convection zone tube flux is also critical to the behaviour of the eruption. The flux rope rise speed, flux rope size, and amount of reconnection between emerging field and coronal field all depend on the strength of the dipole field relative to the strength of the flux tube field.  Very soon after the internal reconnection rate increases, the new coronal flux rope accelerates into the corona, reaching speeds of the order of 60 km/s, which is consistent with the simulations of \\citet{Archontis_Hood_2012}. Further evolution is restricted by the size of the simulation domain. \\citet{Patsourakos_2013} recently observed the formation and eruption of a coronal flux rope, and found that a confined flare occurred immediately prior to the identification of the formation of a flux rope. This rope then rose upwards with speeds of 60 km/s for about 7 hours, before an eruptive flare occurred, with the complete eruption of the coronal flux rope. {\\color{blue} These observations suggest that the simulations in this paper are only capturing the formation and early eruption of a coronal flux rope. Furthermore, the relatively small domain and fast timescales in these simulations, when compared to an active region CME, mean that further studies with larger domains and active regions are required to see if dynamic flux emergence can capture all the stages of a CME eruption, particularly the higher eruption speeds observed in studies such as \\citet{Patsourakos_2013}. }  {\\color{blue} We note that the formation of an eruptive flux rope in these simulations is markedly different from that in the simulation of \\citet{Roussev_2012}, where a similar geometry is used (emerging magnetic field interacting and reconnecting with a coronal dipole). In the simulations of \\citet{Roussev_2012} a coronal flux rope is formed by reconnection between the emerging field and the coronal dipole (not due to internal reconnection as in the simulations in this paper). In \\citet{Roussev_2012} null points are formed at either side of the emerging field region, and a plasmoid is formed at the apex of the emerging region, with flux from both dipole and emerging structure adding to this plasmoid, as can be seen in Figure 1 of \\citet{Roussev_2012}. Because the reconnection which forms the flux rope in \\citet{Roussev_2012} occurs in a region where little of the axial field of the convection zone flux tube has emerged, the resulting flux rope is highly twisted (with $\\sim 10$ windings observed). The flux rope in the simulations in this paper is formed by internal reconnection (the external reconnection with the dipole field aids the emergence of the sub-surface field into the corona). Thus the reconnection occurs in a region where a significant portion of the axial field of the convection zone flux tube has emerged, and so the resulting flux rope has relatively less twist (typically only two windings in the corona are seen). These differences between the simulations of \\citet{Roussev_2012} and those of this paper are most likely caused by differences in how the flux emergence is treated. In \\citet{Roussev_2012}, the emergence at the lower coronal boundary is driven by surface data from a separate flux emergence simulation which does not include a coronal field, and this data is spatially scaled and modified in magnitude to fit coronal conditions. In our simulations the flux emergence and its effect in the corona are both solved self-consistently in a single computational domain. }   These simulations exhibit eruptive behavior of coronal flux ropes \\textit{immediately} after formation, which adds some evidence to the claim that flux ropes are formed during an eruption, not prior to it. However, the above argument suggests we may be covering only a portion of a typical active region CME, and so definitive statements on this matter are difficult until further simulations are performed.  One drawback of successful CME models such as the magnetic breakout model \\citep[e.g.][]{2008ApJ...683.1192L,2008ApJ...680..740D} and the flux rope CME models which rely on the torus instability \\citep{2003ApJ...588L..45R,2005ApJ...630L..97T,2008ApJ...684.1448M} is that they do not dynamically emerge the sheared magnetic field that is required to initiate the eruption from its origins in the convection zone. Typical examples of how this emergence is modeled are via shearing and rotational flows applied to the model surface, which creates sheared field and initiates eruptions (in the case of the breakout model), or by assuming the coronal flux rope is pre-formed in the corona, or by forming a flux rope via magnetic flux cancellation at the surface (in the case of the CME models which rely on the torus instability). These approaches separate the CME models from the source from which they derive their magnetic energy: the convection zone and the solar dynamo. To better predict space weather, one important step is to eliminate this separation. The simulations shown in this paper, which include simple improvements on previous works, such as the use of a dipole field in the corona, are able to self-consistently create many of the surface and coronal signatures used by some CME models. These signatures include surface shearing and rotational flows, quadrupolar geometry above the surface, central sheared arcades reconnecting with oppositely orientated overlying dipole fields, the formation of coronal flux ropes, and internal reconnection which resembles the classical flare reconnection scenario. Thus, within these simulations we have validated the use of  these proxies for flux emergence by certain CME models, and made a major step forward towards fully self-consistent models of the buildup and eruption of magnetic energy in the corona.%"
    },
    "1402/1402.6470_arXiv.txt": {
        "abstract": "We analyze morphologies of the host galaxies of 35 X-ray selected active galactic nucleus (AGNs) at $z\\sim2$ in the Cosmic Evolution Survey (COSMOS) field using Hubble Space Telescope/WFC3 imaging taken from the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS). We build a control sample of 350 galaxies in total, by selecting ten non-active galaxies drawn from the same field with the similar stellar mass and redshift for each AGN host.  By performing two dimensional fitting with GALFIT on the surface brightness profile, we find that the distribution of \\ser index (n) of AGN hosts does not show a statistical difference from that of the control sample.  We measure the nonparametric morphological parameters (the asymmetry index A, the Gini coefficient G, the concentration index C and the $M_{20}$ index) based on point source subtracted images. All the distributions of these morphological parameters of AGN hosts are consistent with those of the control sample. We finally investigate the fraction of distorted morphologies in both samples by visual classification. Only $\\sim 15\\%$ of the AGN hosts have highly distorted morphologies, possibly due to a major merger or interaction. We find there is no significant difference in the distortion fractions between the AGN host sample and control sample. We conclude that the morphologies of X-ray selected AGN hosts are similar to those of nonactive galaxies and most AGN activity is not triggered by major merger. ",
        "introduction": "There is a long-standing question about what triggers AGN activity. Even up till now, fifty years after quasars (QSOs, i.e. luminous AGNs) have been discovered (Schmidt 1963), there is no convincing answer yet. Major merger is a preferred candidate under the framework of hierarchical structure formation (e.g. Hopkins et al. 2006). Other secular mechanisms, such as minor merger, disk/bar instabilities, colliding clouds, supernova explosions, have also been proposed to fuel black hole accretion and trigger AGN activity (see Kormendy \\& Kennicutt 2004; Martini 2004; Jogee 2006 for reviews).  One possible approach is to analyze morphologies of AGN host galaxies and compare them with those of non-active galaxies. However, most morphological analysis of AGN hosts can be seriously biased if AGNs contribute significantly the total flux (e.g. Gabor et al. 2009; Pierce et al. 2010; B$\\ddot{o}$hm et al. 2013). A high-resolution image and a careful decomposition of a point-like source and AGN host will be necessary. Until recently, with the advent of Hubble Space Telescope (HST), such morphological analyses for large AGN samples at $z<1.3$ and $z\\sim2$ in the rest-frame optical have been taken with Advanced Camera for Surveys (ACS) and Wide Field Camera 3 (WFC3), respectively.  Previous studies have found different results. At low redshift ($z<1.3$), some works found a significant fraction of AGNs have bulge-dominated morphologies from a non-parametric classification (e.g. Grogin et al. 2005; Pierce et al. 2007), which could be biased by the point-source component which is not removed in these studies. While other works with a careful point-source removal found that AGN hosts show a wide range of morphologies between bulge and disk dominated (Gabor et al. 2009) or over half of them have a significant disk component (Cisternas et al. 2011). In these studies, no enhancement of merger or interaction signatures have been found in AGN hosts by comparing them with control samples. Urrutia et al. (2008), however, found a high merger fraction (11 out of 13 in their luminous red QSOs sample). This may lead to an explanation that merger fraction is dependent on AGN luminosity (Treister et al. 2012, but see also Villforth et al. 2013). At $z\\sim2$, AGN hosts  seem to be disk-dominated with a similar merger fraction as those of nonactive galaxies (Schawinski et al. 2011,2012; Simmons et al. 2012) or have a significant fraction of disk galaxies with bulge component (Kocevski et al. 2012).  In this Letter we take the comprehensive structural and morphological analyses of X-ray selected AGN host galaxies in the CANDELS-COSMOS field using HST/WFC3 H-band imaging  at $1<z<3$. We try to find whether there is a special morphological type which AGN hosts prefer to belong to. And we examine whether there is a significant difference of % morphological properties between AGN hosts and control samples. In section 2, we describe our sample selections of X-ray selected AGN and control samples. In section 3, we describe the methods to measure the morphological parameters. We at first perform two dimensional light profile fitting with GALFIT using a \\ser model representing host galaxy contribution and a PSF model representing the point source contribution. Then based on the point source subtracted images, we compute the nonparametric morphological parameters. Finally we perform the visual classification based on the H-band image and the residual image. In section 4, we give the main results of our morphological analyses. In section 5, we summarize and discuss our results.  Throughout this letter, we assume a concordance $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, $H_{\\rm 0}=70$ $\\kms$ Mpc$^{-1}$. ",
        "conclusions": "In this letter, we analyze the structure and morphology of X-ray selected AGN host galaxies in the CANDELS-COSMOS field using HST/WFC3 imaging in H band at $z\\sim2$ and compare them with those of a mass-matched non-active galaxy sample. Our primary findings are as follows: 1) Near two thirds of X-ray selected AGN host galaxies in our sample at $z\\sim2$ have disk-like morphologies (including disk-dominated and disk with a bulge component morphologies), while a significant fraction (over one thirds) of them are bulge dominated; 2) All structural (\\ser index n and effective radius $R_e$) and morphological (Gini, $M_{20}$,concentration and asymmetry) parameters of AGN host galaxies in our sample have the similar distributions as those of non-active galaxies. We conclude that, from the point of view of structure and morphology, AGN host galaxies and non-active galaxies are indistinguishable. From other point of view of such as color, color gradients and stellar population properties, Rosario et al (2013) also found that X-ray selected AGN host galaxies and non-active galaxies are indistinguishable. 3) Only a small fraction ($\\sim 15\\%$) of X-ray selected AGN host galaxies have major merger signatures according to our visual classification. Compared to the mass-matched non-active galaxy sample, AGN hosts do not show a significant excess of distorted morphologies.  As shown by previous host galaxy simulations (e.g., Simmons \\& Urry 2008; Gabor et al. 2009; Pierce et al. 2010), the result that AGN host galaxies are disk-dominated characterized by a low fitted \\ser index n is reliable. While disk contribution cannot be ruled out in bulge-dominated galaxy with a high fitted \\ser index n. Simmons \\& Urry (2008) found that bulge-dominated AGN host galaxies  with $n>4$ may have a significant disk contribution (up to 45\\% of their total host galaxy light). The disk fraction ($\\sim63\\%$) of AGN host galaxies in our sample should be taken as a lower limit. This value is a bit lower than the disk fraction $75\\%$ in Simmons et al. (2012), and is much lower than the disk fraction $\\sim90\\%$ in Schawinski et al. (2011). The relatively lower disk fraction in our sample could be the result of different sample selections. The X-ray luminosities of AGNs in our sample are on average one order of magnitude higher than those in Schawinski et al. (2011).  Complemented with several previous findings at $0<z<3$ (e.g., Gabor et al. 2009, Cisternas et al. 2011, Schawinski et al. 2011, Simmons et al. 2012, Kocevski et al. 2012), some consistent results can be summarized for X-ray selected AGNs with $10^{42}$ erg s$^{-1}$ $<L_X<10^{45}$ erg s$^{-1}$  up to $z\\sim3$: disk morphologies ; no enhanced distorted morphologies and similar structure and morphology as non-active galaxies. These indicate that major mergers are not necessary for triggering AGN activities in the galaxies of the X-ray selected sample. Secular internal processes, such as gravitational instabilities  and dynamical friction, should play a crucial role in triggering X-ray selected AGN activities and black hole-host galaxy co-evolution. However, our findings do not conflict with supermassive black hole (SMBHs)-host galaxy co-evolution  scenarios in which major mergers are responsible for triggering both star formation and luminous QSO activities (e.g., Hopkins et al. 2006).  Unlike X-ray selected AGNs, luminous QSOs commonly have very high fraction (up to $\\sim 100\\%$) in major merger and are specifically in dust-shrouded, merger-induced starbursts (e.g., Urrutia et al. 2008; Treister et al. 2012). Thus two different modes of BH triggering and growth have been indicated: the BH growth by secular evolution of disk galaxies is important for low-luminosity AGNs, while major mergers may feed luminous quasars (e.g., Kormendy \\& Ho 2013)."
    },
    "1402/1402.4828_arXiv.txt": {
        "abstract": "By means of a semi-analytic model of galaxy formation, we show how the local observed relation between age and galactic stellar mass is affected by assuming a DM power spectrum with a small-scale cutoff.  We compare results obtained by means of both a $\\Lambda$-cold dark matter ($\\Lambda$CDM) and a $\\Lambda$-warm dark matter ($\\Lambda$WDM) power spectrum - suppressed with respect to the $\\Lambda$CDM at scales below $\\sim 1$ Mpc. We show that, within a LWDM cosmology with a thermal relic particle mass of 0.75 keV, both the mass-weighted and the luminosity-weighted age-mass relations are steeper than those obtained within a $\\Lambda$CDM universe, in better agreement with the observed relations. Moreover, both the observed differential and cumulative age distributions are better reproduced within a $\\Lambda$WDM cosmology. In such a scenario, star formation appears globally delayed with respect to the $\\Lambda$CDM, in particular in low-mass galaxies. The difficulty of obtaining a full agreement between model results and observations is to be ascribed to our present poor understanding of baryonic physics. ",
        "introduction": "The current `concordance' cosmological paradigm, which envisages a Universe dominated by cold dark matter (DM) and a dark energy $\\Lambda$, predicts a `bottom-up' formation of structures, in which the smallest DM halos are the first to collapse, then interact and later merge to form larger halos. This $\\Lambda$-cold dark matter ($\\Lambda$CDM ) scenario is able to account for several lines of evidence, which concern primarily the large-scale properties of the Universe  and which include the temperature anisotropies in the microwave background radiation (Komatsu et al. 2011), galaxy clustering (Percival et al. 2010) and the observed flux power spectrum of the Lyman-$\\alpha$ forest (Viel et al. 2005).  This picture motivated several theoretical investigations of the properties of galaxies, which soon evidenced some major problems in the dominant cosmological picture, in particular on galactic and sub-galactic scales, i.e  below a few Mpc. First of all, N-body simulations have shown that a $\\Lambda$CDM  cosmology implies inner density profiles much steeper than those inferred from the rotation curves of real galaxies (Moore et al. 1999, Gentile et al. 2007). In addition, $\\Lambda$CDM-based simulations overpredict the number of satellites residing in  Milky Way-sized halos by $5-10$ times (Kauffmann et al. 1993; Moore et al. 1999, Klypin et al. 1999), a discrepancy commonly referred to as the Missing Satellite problem.  It is worth stressing that other explanations could be found for the missing satellite problem. On the observational side, this problem could be due to the inability of modern telescopes to detect the faintest galaxies (Bullock et al. 2010), or to an overestimate of the virial mass of the Milky-Way dark matter halo (Wang et al. 2012). Moreover, it was shown that various baryonic processes, as well as large-scale effects such as `cosmic web stripping' (Benitez-Lambay et al. 2013) can play some role in reducing the abundance of low-mass objects in MW-sized halos (Sawala et al. 2013). However, it was also shown that the tension between the dynamical properties of Milky-Way satellites and $\\Lambda$CDM expectations are more difficult to reconcile by means of baryonic processes (Boylan-Kolchin, Bullock \\& Kaplinghat 2011; 2012).  Other problems related to the observable properties of stellar populations have been evidenced by $\\Lambda$CDM-based semi-analytical galaxy formation models, which allow one to link directly baryon-related physical processes such as star formation, cooling and feedback (Cole et al. 2000) to the merging histories of DM haloes. In the past decades, semi-analytical models (SAMs) have turned out as useful tools to obtain a physically-motivated statistical view of the global properties of galaxies (Baugh 2006; Silk \\& Mamon 2012). SAMs are commonly used to investigate the main galactic observables related to their star formation histories, their stellar mass and halo mass assembly, including the evolution of the luminosity function (e. g., Somerville et al. 2012, Menci et al. 2012), the cosmic star formation history (e. g., Nagashima et al. 2004), as well as the basic scaling relations (e. g., Fontanot et al. 2009) and the evolution of their chemical abundances (e. g., Calura \\& Menci 2009; Yates et al. 2013), often outlining further problems in reproducing the observations within $\\Lambda$CDM models, again mostly related to the small-scale power excess which manifests in an excess of of faint systems at any redshift.  Although it is still possible that the failure of an accurate description of the statistical properties of galaxies may be due to an incomplete theoretical understanding of the basic physics regulating the most fundamental baryonic processes (see, e.g., Koposov et al. 2009, Kang et al. 2013, Brooks et al. 2013), the above problems could also be related to the intimate nature of dark matter. In fact, the intrinsic excess of substructure typical of the $\\Lambda$CDM scenario is less pronounced in a Universe dominated by Warm Dark Matter (WDM). Compared to CDM, the WDM constituents are lighter particles with masses of a few keV, and hence with larger thermal speeds. In terms of power spectrum, this corresponds to a suppression of the growth of structures on the smallest scales, hence providing a natural mechanism to limit the wealth of substructures typical of the CDM scenario.  The effects of a WDM power spectrum on the statistical properties of galaxies have been investigated in Menci et al. (2012), where it was shown that the adoption of a WDM scenario helps alleviate the tension between model predictions and the most small-scale sensitive observables, such as the faint end of the luminosity function. Also, the evolution of the active galactic nucleus (AGN) luminosity function (LF),  suffering the same problems as the galactic LF, appears better described with the assumption of an underlying WDM power spectrum (Menci et al. 2013).  So far, little attention has been paid to another long-standing problem related to the properties of baryons, i.e. the inability of $\\Lambda$CDM models to account for the observed relation between the luminosity-weighted age of the stellar populations and the galactic stellar mass. The observational study by Gallazzi et al. (2005) indicates that, on average, the larger the galaxy, the older its stellar populations, supporting the well-established galactic `downsizing' (Cowie et al. 1996; Mateus et al. 2006; Fontanot et al. 2009). In principle, this result is not inconsistent with the CDM scenario; in fact, the progenitor of massive galaxies are predicted to form in biased, high-density regions of the density field where star formation starts at early epochs. At later cosmological epochs, bottom-up hierarchical structure formation then assembles these early-formed stellar populations into massive, low-redshift galaxies. However, $\\Lambda$CDM  models also predict the early collapse of a huge number of low-mass halos, which remain  isolated at later times retaining the early-formed stellar populations; as a result, $\\Lambda$CDM-based SAMs  generally provide flat age-mass relations (Fontanot et al. 2009; Pasquali et al. 2010; De Lucia \\& Borgani 2012), indicating a basic coevality of satellites and central galaxies. The suppression of the small-scale power due to a different free streaming length is expected to have effects not only on the abundance of collapsed objects below a cut-off scale, but also on the time of formation of halos (Angulo et al. 2013), consequently of the formation and assembly history of baryonic structures and ultimately on the ages of stellar populations. In this paper, our aim is to study the effects of a WDM cosmology on the star formation of galaxies. In particular, we will address the issue of the ages of stellar populations in galaxies of various masses within a WDM-based semi-analytical model (SAM) for galaxy formation, and we will test whether such a scenario can help alleviate the above discrepancy. \\\\ This paper is organized as follows. In Sect. 2, we briefly describe the main ingredients of the SAM used in this work. Our results are presented in Sect. 3. Finally, our conclusions are drawn in Sect. 4. ",
        "conclusions": "At galactic and sub-galactic scales ($\\lesssim 1 Mpc$), computations based on N-body simulations pose several problems to the $\\Lambda$CDM scenario. The first problem concerns the number of satellites present in the Milky Way halo, which in  $\\Lambda$CDM simulations is expected to contain between $\\sim 5$  and $\\sim 10$  more satellites than what current observations indicate.  On the observational side, the Missing Satellite problem could also be related to the limited capability of modern telescopes to detect the lowest mass galaxies: the number of detected Mily-Way satellites is increasing all the time as we probe fainter objects, and 5-20 times more could be discovered in the future (Bullock et al. 2010). On the other hand, a theoretical explanation for the Missing Satellite problem could also be found within the $\\Lambda$CDM framework. Wang et al. (2012) have shown that the overabundance of local satellites could be related to our poor knowledge of the virial mass of the Milky Way halo, and that with a dark matter halo of mass $M_h < 10^{12} M_{\\odot}$, the discrepancy between simulation results and observations can be reconciled. Sawala et al. (2013) have shown instead that sub-galactic baryonic processes could alleviate this tension between simulation results, as could large-scale effects such as 'cosmic stripping' (Benitez-Llambay et al. 2013). There are however other problems which seem more difficult to solve within $\\Lambda$CDM. Boylan-Kolchin, Bullock \\& Kaplinghat (2012) showed that the bright satellites of the MW are hosted by DM sub-halos significantly less dense than the most massive sub-halos in state-of-the-art $\\Lambda$CDM simulations. (Boylan-Kolchin, Bullock \\& Kaplinghat 2012). In this case, even strong baryonic feedback does not seem to alleviate this discrepancy. The origin of the above problems could be related to an excess of power of the CDM power spectrum at small scales; in this case, a natural solution is represented by the adoption of a universe dominated by warm dark matter.  In a previous paper, Menci et al. (2012) have tested for the first time the effects of $\\Lambda$WDM on the statistical properties of galaxies at various redshifts. These authors have shown that the adoption of a power spectrum  which is suppressed (compared to the standard CDM case) below a cut-off scale of $\\sim $ 1 Mpc, helps limit the wealth of substructure typical of LCDM models.  By means of the same model described in Menci et al. (2012), in this paper we have studied the effects of a $\\Lambda$WDM power spectrum on the ages of the stellar populations and on the star formation histories of present-day galaxies. We have focused on the age-mass relation, a well-established property of local galaxies which outlines one fundamental aspect of their `downsizing' nature, i.e. the fact that the largest galaxies host the oldest stellar populations (Gallazzi et al. 2005; Fontanot et al. 2009; Pasquali et al. 2010). Explaining this relation within `concordance' $\\Lambda$CDM cosmological models has represented a long-standing problem, for the main reason that in the standard cosmological scenario, local low-mass galaxies appear too passive, characterised by too low present star formation rates and by an excess of old stars, with respect to the observations. This manifests in a predicted age-mass relation which is too shallow with respect to the observed one (Fontanot et al. 2009; De Lucia \\& Borgani 2012), with a reasonable agreement between models and data at the bright end, but accompanied by a severe overestimate of the stellar ages of galaxies at masses $M_{*}<10^{10} M_{\\odot}$. This discrepancy is partly related to the well-known difficulties of modelling baryonic processes in cosmological models: star formation, gas cooling and energetic feedback from exploding supernovae are highly non-linear processes and our present understanding of each of them is far from being complete. In this regard, a step forward is represent by the study of Henriques et al. (2013) who showed how, by means of a new physical mechanisms allowing to reincorporate the gas ejected in the lowest mass galaxies, it was possible to obtain local dwarf galaxies bluer and more actively star forming than in previous studies.  However, the inability of $\\Lambda$CDM models to describe realistically the star formation histories of low-mass galaxies could be due to the underlying $\\Lambda$CDM cosmological framework, in which the smallest structures are the earliest to emerge and, without any realistic physical mechanism which delays the baryon cooling and the subsequent star formation, these objects form too large a fraction of their stars (compared to the observational estimate) at high redshift (Fontanot et al. 2009). In this regard, the choice of a WDM power spectrum may help alleviating several discrepancies between the observed properties of galaxies and the model results since, beside implying the suppression of the small scale power, it also causes an overall delay of the growth of structure formation (Amendola et al. 2013; Destri et al. 2013) and of star formation. Note however that, in order to enlighten the impact of a  WDM power spectrum on the age distributions, we have adopted in this paper a particle mass corresponding (for thermal relics) to $m_X=0.75$ keV. Several observational constraints indicate larger values for the dark matter particle mass: from long GRBs at high redshifts, de Souza et al. 2013 suggest $m_X>1.6$ keV at 95 per cent confidence level. From the analysis of image fluxes in gravitationally lensed quasi-stellar objects, Miranda \\& Macci{\\`o} (2007) find a lower limit of 10 keV for the mass of WDM candidates in the form of sterile neutrinos. More constraints come from various recent studies based on the extensive use of N-body simulations. From the distribution of satellites in the MW, Macci{\\`o} \\& Fontanot (2010) suggest a lower limit of 1 keV for thermal WDM particles. The latest constraints from the observed absorption in the Lyman-$\\alpha$ forest indicate a lower limit of $m_X\\gtrsim 3.3$ keV for thermal relic particles (Viel et al. 2013; Markovic \\& Viel 2013). From high-redshift galaxy counts, Schultz et al. (2014) suggest that a 1.3 keV WDM particle can be ruled out at high significance. Overall, most of these results would indeed leave very little room for a contribution of WDM to the solution of the small scale crisis of CDM (see, e.g., Polisensky \\& Ricotti 2013; Herpich et al. 2013). However, it is worth stressing that artificial fragmentation is a serious issue in current N-body numerical simulations on scales around and below the free streaming length (Lovell et al. 2012, Schneider et al. 2013, Hahn et al. 2013; Markovic \\& Viel 2013; Lovell et al. 2013), and that increasing the resolution even by rather large factors is not sufficient to prevent this effect (Wang \\& White 2007). The recent study of Lovell et al. (2013) uses an algorithm to identify spurious clumps and to remove them from the catalogues and, from the constrains provided by the abundance of Milky-Way satellites, they obtain a conservative lower limit to the WDM thermal particle mass $m_X$=1.5 keV. The use of a similar algorithm also in other studies based on N-body simulations will be crucial to investigate further the systematic effects of artificial fragmentation, and to eventually have more robust constraints on the WDM particle mass candidate.  In such a context, the scope of the present paper is not to provide a test for the WDM scenario, nor to pin down the DM candidate that provides the power spectrum that best matches the observations. Rather, we aim at showing that {\\it assuming} that the solution of the CDM small scale crisis relies in a cut-off in the power spectrum, such a cut-off {\\it also} provides a viable solution to the problems related to the age distribution of stellar populations. In particular, the assumption of a DM power-spectrum cutoff such as that provided by a WDM particle with mass $m_X\\approx 1$ keV has a significant impact on the age distribution of stellar populations in local galaxies as detailed below:  \\begin{itemize} \\item Within a $\\Lambda$WDM cosmology, both the mass-weighted and the luminosity-weighted age-mass relation are steeper than those obtained within a $\\Lambda$CDM universe, in better agreement with the one estimated observationally. As expected, $\\Lambda$WDM galaxies of all masses  present, on average,  younger stellar populations than their $\\Lambda$CDM counterparts. At the massive end, $\\Lambda$WDM galaxies appear younger than real galaxies. \\item Our study of the average age distribution of local galaxies has confirmed that the best results can be obtained within a LWDM cosmology, assuming a particle mass of 0.75 keV: in this case, the bimodal behaviour of the observational distribution is reproduced, although differences in the width and in the peak ages between the predicted distribution and the observed one are visible. \\item We studied the effects of a WDM power spectrum on the average star formation histories of the simulated galaxies. We grouped the present-day stellar masses of the model galaxies into four different mass bins and in each bin, we computed an average star formation rate. The adoption of a $\\Lambda$WDM cosmology affects mostly low-mass galaxies, for which we find a star formation history significantly different with respect to the $\\Lambda$CDM. While in the $\\Lambda$CDM the SFH peaks at the earliest times and then decreases progressively, in the $\\Lambda$WDM case the behaviour is reversed, in that the average SFH increases with cosmic time. This result seems in qualitative agreement with recent attempts to reconstruct the SFHs of blue galaxies by means of physically motivated, comprehensive libraries of synthetic spectra (Pacifici et al. 2013). Significant differences are visible also in the highest mass bin ($11<log(M/M_{\\odot})$), where galaxies within $\\Lambda$WDM show a lower early peak and larger present-day SFRs than their $\\Lambda$CDM counterparts. This aspect will need to be improved in the future, by testing possible mechanisms to quicken the gas consumption timescales in the most massive systems. \\end{itemize} Other studies carried out by means of semi-analytical models have demonstrated the importance of the choice of the parameters related to baryonic physics in determining some fundamental scaling relations and properties of galaxies (Kang et al. 2013) and AGNs (Menci et al. 2013) within WDM models. Kang et al. (2013) outlined the difficulty in disentangling between the true effects due to WDM and the specific implementation of baryonic physics when observables such as luminosity functions, the ratio between  between DM and stellar mass in galaxies and, most of all, dynamical mass estimators such as the Tully-Fisher relation are considered. Menci et al. (2013) studied the evolution of the AGN LF and showed how a $\\Lambda$CDM model with extreme assumptions for both AGN and stellar feedback can produce results very similar to those obtained by means of a WDM, in particular at low redshift. The tests performed in this paper have shown that the stellar populations in galaxies of stellar mass $<10^{10}$ M$_{\\odot}$, characterised by a relatively simpler physics owing mostly to the minor effects of AGNs on their star formation histories, may help disentangling between the effects of a WDM power spectrum and the baryonic physics. We have shown that a WDM cosmology can naturally account for younger stellar populations in low mass galaxies with respect to a $\\Lambda$CDM model. At present, figuring realistic mechanisms which can further delay star formation in low-mass systems at high redshift seems difficult and may require substantial revisions in the current feedback schemes (Fontanot et al. 2009, Weinmann et al. 2012, Henriques et al. 2013).   \\begin{figure*} \\centering \\includegraphics[width=150mm]{sfh.pdf} \\caption{Average star formation histories of our simulated galaxies divided into four stellar mass bins (reported on  top of each panel), computed adopting a $\\Lambda$WDM cosmology (dashed, red curves) and a $\\Lambda$CDM cosmology (blue solid curves).} \\label{fig_sfh} \\end{figure*}"
    },
    "1402/1402.4705_arXiv.txt": {
        "abstract": "We have compiled photometry at 3.4, 4.6, 12 and 22 $\\mu$m from the all-sky survey performed by the {\\it Wide-field Infrared Survey Explorer} ({\\it WISE}) for all known members of the Taurus complex of dark clouds. Using these data and photometry from the {\\it Spitzer Space Telescope}, we have identified members with infrared excess emission from circumstellar disks and have estimated the evolutionary stages of the detected disks, which include 31 new full disks and 16 new candidate transitional, evolved, evolved transitional, and debris disks. We have also used the \\wise All-Sky Source Catalog to search for new disk-bearing members of Taurus based on their red infrared colors. Through optical and near-infrared spectroscopy, we have confirmed 26 new members with spectral types of M1 -- M7. The census of disk-bearing stars in Taurus should now be largely complete for spectral types earlier than $\\sim$M8 ($M \\gtrsim 0.03$~$M_\\odot$). ",
        "introduction": "Complete samples of circumstellar disks in star-forming regions and accurate classifications of those disks represent a foundation for studies of star and planet formation. Both the identification of circumstellar disks and their classification are most easily performed using mid-infrared (IR) continuum emission from warm circumstellar dust. Because the atmosphere is bright and strongly absorbing at mid-IR wavelengths, space-based telescopes have provided the most sensitive data of this kind. The all-sky mid-IR images from the {\\it Infrared Astronomical Satellite} \\citetext{IRAS; \\citealt{neu84}} enabled the first detections of circumstellar disks in star-forming regions. The most thorough census of disks was performed in Taurus because it is nearby ($d=140$~pc; \\citealt{wic98,loi05,tor07,tor09}) and has a low enough stellar density that its members could be resolved by {\\it IRAS} \\citep{ken95}. Subsequent mid-IR telescopes, such as the {\\it Infrared Space Observatory} \\citep{kes96} and the {\\it Spitzer Space Telescope} \\citep{wer04}, have offered progressively better sensitivity and resolution, enabling detections of disks at fainter levels and in more crowded fields. Due to its modest field of view (5$\\arcmin$), \\spitz primarily observed more compact star-forming regions (\\citealt{lad06}; \\citealt{sic06}; \\citealt{dah07}; \\citealt{her07}; \\citealt{luh08}; \\citealt{gut09}), although it was able to map a significant fraction of a few widely distributed populations such as Taurus (\\citealt{luh10tau}; \\citealt{reb10}).   The latest mid-IR satellite, the {\\it Wide-field Infrared Survey Explorer} ({\\it WISE}; \\citealt{wri10}), has lower spatial resolution than \\spitz but it covered the entire sky. As a result, \\wise provides mid-IR photometry for the portions of large star-forming regions and associations that were not imaged by {\\it Spitzer}. For instance, \\wise data have been used to search for new disk-bearing stars in Taurus \\citep{reb11} and to classify disks among the known members of Upper Sco \\citep{luh12}. Because of the importance of Taurus in studies of disks and because proper use of \\wise data requires great care, we have performed our own search for new members with disks in Taurus, and have classified all the disks that we have detected around the known members. In this paper we begin by compiling photometry from 3 to 24 $\\mu m$ from both {\\it WISE} and \\spitz for all known members of Taurus (Section \\ref{sec:photo}).  With these data, we then measure mid-IR excesses (Section \\ref{sec:excess}) and classify the evolutionary stages of the detected disks (Section \\ref{sec:class}).  Finally, we use \\wise photometry in conjunction with proper motions, color-magnitude diagrams, and follow up spectroscopy to identify new members of Taurus (Section \\ref{sec:search}). ",
        "conclusions": "We have performed a survey of circumstellar disks in the Taurus star-forming region using {\\it WISE}. Our results are summarized as follows:  \\begin{enumerate}  \\item We have examined all images from the cryogenic phase of \\wise for all known members of Taurus to check  for false detections, blends with nearby objects, and extended emission. We have presented the resulting catalog of vetted photometry in the \\wise bands at 3.4, 4.6, 12, and 22 $\\mu$m ($W1$--$W4$). All resolved, unblended members are detected by \\wise in at least one band.  \\item By using colors constructed from $K_s$ and six \\spitz and \\wise bands, we have identified Taurus members showing excess emission from circumstellar disks and have estimated the evolutionary stages of the detected disks, consisting of full, transitional, evolved, evolved transitional, and debris disks. Our classifications are generally consistent with those found by \\cite{luh10tau} and \\cite{fur11}. We have found 31 new full disks and 16 new candidate disks in the more advanced evolutionary stages.  \\item Using photometry from {\\it WISE}, {\\it Spitzer}, 2MASS, and USNO, proper motions from UCAC4 and other catalogs, and our optical and near-IR spectroscopy, we have found 26 new members of Taurus with spectral types of M1 -- M7. The census of disk-bearing stars in Taurus should now be largely complete for $\\mw1<13$ ($\\lesssim$M8;  $M \\gtrsim 0.03$~$M_\\odot$).   \\end{enumerate}"
    },
    "1402/1402.2779_arXiv.txt": {
        "abstract": "We present an analysis of multi-timescale variability in line-of-sight X-ray absorbing gas as a function of optical classification in a large sample of Seyfert active galactic nuclei (AGN) to derive the first X-ray statistical constraints for clumpy-torus models. We systematically search for discrete absorption events in the vast archive of \\textit{Rossi X-ray timing Explorer} monitoring of dozens of nearby type I and Compton-thin type II AGN. We are sensitive to discrete absorption events due to clouds of full-covering, neutral or mildly ionized gas with columns $\\ga 10^{22-25}$~cm$^{-2}$ transiting the line of sight.  We detect 12 eclipse events in 8 objects, roughly tripling the number previously published from this archive. Peak column densities span $\\sim4-26 \\times 10^{22}$~cm$^{-2}$, i.e., there are no full-covering Compton-thick events in our sample. Event durations span hours to months. The column density profile for an eclipsing cloud in NGC~3783 is doubly spiked, possibly indicating a cloud that is being tidally sheared.  We infer the clouds' distances from the black hole to span $\\sim 0.3 - 140 \\times 10^{4} R_{\\rm g}$. In seven objects, the clouds' distances are commensurate with the outer portions of Broad Line Regions (BLR), or outside the BLR by factors up to $\\sim10$ (the inner regions of infrared-emitting dusty tori). We discuss implications for cloud distributions in the context of clumpy-torus models. Eight monitored type II AGN show X-ray absorption that is consistent with being constant over timescales from 0.6 to 8.4~yr. This can either be explained by a homogeneous medium, or by X-ray-absorbing clouds that each have $N_{\\rm H} \\ll 10^{22}$~cm$^{-2}$. The probability of observing a source undergoing an absorption event, independent of constant absorption due to non-clumpy material, is $0.006^{+0.160}_{-0.003}$ for type Is and $0.110^{+0.461}_{-0.071}$ for type IIs. ",
        "introduction": "While it is generally accepted that active galactic nuclei (AGN) are powered by accretion of gas on to supermassive ($10^{6-9}\\Msun$) black holes, the exact geometry and mechanisms by which material gets funneled from radii of kpc/hundreds of pc down to the accretion disk at sub-pc scales remains unclear. This unsolved question is, however, linked to the overall efficiency of how supermassive black holes are fed and to AGN duty cycles.  Clues to the morphology of the circumnuclear gas come from different studies across the electromagnetic spectrum. For instance, Seyfert AGN are generally classified into two broad categories based on the detection or lack of broad optical/UV emission lines (type I or II, respectively), with intermediate subclasses (1.2, 1.5, 1.8, 1.9) depending on broad lines' strengths (e.g., Osterbrock 1981). Unification theory posits that all Seyferts host the same central engine, but observed properties depend on orientation due to the presence of dusty circumnuclear gas blocking the line of sight to the central illuminating source in type II Seyferts.  The classical model holds that a primary component of this obscuring gas is a pc-scale, equatorial, dusty torus (Antonucci \\& Miller 1985; Urry \\& Padovani 1995), where ``torus'' generally denotes a donut-shaped morphology, supplying the accretion disk lying co-aligned inside it. However, some recent papers (e.g., Pott et al.\\ 2010) use the term ``torus'' to denote simply the region where circumnuclear gas can exist, with the precise morphology of that region still to be determined; we adopt that notation in this paper.  A donut morphology can explain the observed bi-conical morphology of Narrow Line Regions (NLR), e.g., via collimation of ionizing radiation (e.g., Evans \\etal\\ 1994), and also why most type II Seyferts show evidence for X-ray obscuration along the line of sight.  However, the relation between optical/UV-reddening dust and X-ray absorbing gas is not straightforward.  Very roughly 5 per cent of all AGN have differing X-ray and optical obscuration indicators (e.g., X-ray-absorbed type Is, Perola \\etal\\ 2004; Garcet \\etal\\ 2007); a complication is that optical obscuration probes only dusty gas, while X-ray absorption probes both dusty and dust-free gas, and X-ray columns can exceed inferred optical-reddening columns by factors of 3--100 (Maiolino \\etal\\ 2001).  Lutz et al.\\ (2004) and Horst et al.\\ (2006) demonstrated that the infrared (IR) emission by dust, which thermally re-radiates higher energy radiation, is isotropic, with type I and II Seyferts having similar ratios of X-ray to mid-IR luminosity. This is not expected from classical unification models, since the classical donut-shaped torus both absorbs and re-emits anisotropically.  Furthermore, if the torus is comprised of primarily Compton-thick gas, i.e., with X-ray-absorption column densities $N_{\\rm H}$ above $1.5 \\times 10^{24}$~cm$^{-2}$, then it is not clear how Compton-thin columns can be attained unless lines of sight happen to ``graze'' the outer edges.  A unification scheme based solely on inclination angle and the presence of a donut-shaped absorbing morphology may thus be an oversimplification (Elvis 2012).  It is also not clear how geometrically thick structures (scale heights $H/R \\sim 1$) comprised of cold gas (100 K) are supported vertically over long timescales (Krolik \\& Begelman 1988).  Meanwhile, the community has accumulated observations of variations in the X-ray absorbing column $N_{\\rm H}$ in both (optically classified) type Is and IIs, with timescales of variability ranging from hours to years. For instance, Risaliti, Elvis, \\& Nicastro (2002; hereafter REN02) claim that variations in $N_{\\rm H}$ in a sample of X-ray-bright, Compton-thin and moderately Compton-thick type IIs are ubiquitous, with typical variations up to factors of $\\sim1.5-3$. Current X-ray missions such as \\textit{XMM-Newton}, \\textit{Chandra} and \\textit{Suzaku} have enabled high precision studies of variations in $N_{\\rm H}$ in numerous AGN; major findings include the following:  $\\bullet$ Numerous moderately Compton-thick variations ($\\Delta$$N_{\\rm H} \\sim 10^{23-24}$~cm$^{-2}$) on timescales $\\la 1-2$ d in NGC~1365 (Risaliti \\etal\\ 2005, 2007, 2009a) and ESO 323--G77 (Miniutti \\etal\\ 2014); moderately Compton-thick variations on timescales from days to months in NGC~7582 (Bianchi \\etal\\ 2009).  $\\bullet$ Changes in covering fraction of partial-covering absorbers on timescales from $<$1 d (NGC~4151, Puccetti \\etal\\ 2007; NGC~3516, Turner \\etal\\ 2008; Mkn~766, Risaliti \\etal\\ 2011; SWIFT J2127.4+5654, Sanfrutos \\etal\\ 2013) to months--years in NGC~4151 (de Rosa \\etal\\ 2007 and Markowitz \\etal\\ in preparation, using \\textit{BeppoSAX} and \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}) data, respectively).  $\\bullet$ Time-resolved spectroscopy of full eclipse events (ingress and egress), yielding constraints on clouds' density profiles for long-duration (3--6 months) eclipses in NGC~3227 (Lamer \\etal\\ 2003) and Cen~A (Rivers \\etal\\ 2011b) and for eclipses $\\la$1 d in NGC~1365 (Maiolino \\etal\\ 2010) and SWIFT J2127.4+5654 (Sanfrutos \\etal\\ 2013).  These results suggest that the circumnuclear absorbing gas is \\textit{clumpy}, with non-homogeneous or clumpy absorbers being invoked to explain time-variable X-ray absorption as far back as \\textit{Ariel V} and \\textit{Einstein} observations (Barr \\etal\\ 1977; Holt \\etal\\ 1980).  With the concept of a homogeneous, axisymmetric absorber thus under scrutiny, the community has been developing torus models incorporating clumpy gas, e.g., Elitzur \\& Shlosman (2006) and Nenkova \\etal\\ (2008a, 2008b; see H\\\"{o}nig 2013 for a review), although suggestions that the torus should consist of clouds go as far back as, e.g., Krolik \\& Begelman (1986, 1988).  In the most recent models, total line-of-sight absorption for a given source is quantified as a viewing dependent \\textit{probability} based on the size and locations of clouds, although typically, clouds are preferentially distributed towards the equatorial plane.  The fraction of obscured sources depends on the average values and distributions of such parameters as the average number of clouds lying along a radial path, and the thickness of the cloud distribution.  Clouds are possibly supported vertically by, e.g., radiation pressure (Krolik 2007) or disk winds (Elitzur \\& Shlosman 2006).  Observations so far suggest that clouds are typically on the order of $10^{13-15}$~cm in diameter, with number densities $\\sim 10^{8-11}$ cm$^{-3}$. Inferred distances from the black hole, usually based on constraints from X-ray ionization levels and assumption of Keplerian motion, range from light-days and commensurate with clouds in the Broad Line Region (BLR; e.g., for NGC~1365 and SWIFT J2127.4+5654) to several light-months (Rivers \\etal\\ 2011b for Cen~A) and commensurate with the IR-emitting torus in that object. We emphasize ``commensurate,'' as X-ray absorbers lie along the line of sight, but BLR clouds likely contain components out of the line of sight.  Recent IR interferometric observations have spatially resolved distributions of dust down to radii of tenths of pc (e.g., Kishimoto \\etal\\ 2009, 2011; Tristram \\etal\\ 2009; Pott \\etal\\ 2010).  In addition, there are suggestions from reverberation mapping of the thermal continuum emission in four Seyferts that warm dust gas extends down to $\\sim10-80$ light-days, likely just outside the outer BLR (Suganuma \\etal\\ 2006).  As suggested by Netzer \\& Laor (1993), Elitzur (2007), and Gaskell et al.\\ (2008), the BLR and dusty torus may be part of a common radially extended structure, spanning radii inside and outside the dust sublimation radius $R_{\\rm d}$, respectively, since dust embedded in the gas outside $R_{\\rm d}$ suppresses optical/UV line emission.  Optical obscuration is due to dusty gas, while X-ray obscuration can come from dusty or dust-free gas. X-ray obscuration is thus the only way to probe obscuration inside $R_{\\rm d}$.  In order to gain a more complete picture of the geometry of circumnuclear gas in AGN, however, the community needs to gauge the relevance of clumpy-absorber models over a wide range of length scales, including both inside and outside $R_{\\rm d}$, and to clarify the links between the distributions of dusty gas, X-ray-absorbing gas, and the BLR.  To date, however, observational constraints to limit parameter space in clumpy-torus models has been lacking because there has been no statistical survey so far.  One of our goals for this paper is to derive constraints on clumpy-torus models via variable X-ray-absorbing gas, including estimates of the probability that the line of sight to the AGN is intercepted by a cloud.  We use the vast archive of \\textit{RXTE} multi-timescale light curve monitoring of AGN, as described in $\\S$2.  We search for changes in full- or partial-covering $N_{\\rm H}$ in Seyferts.  We use a combination of light curve hardness ratios and time-resolved spectroscopy to identify and confirm eclipses, which are summarized in $\\S$3. We use the observed eclipse durations and the observation sampling patterns to estimate the probability to observe a source undergoing an eclipse in $\\S$4.  In $\\S$5, we infer the eclipsing clouds' radial locations and physical properties, relate the X-ray-absorbing clouds to other AGN emitting components, and describe the resulting observational constraints for key parameters of clumpy-torus models.  The results are summarized in $\\S$6.  In a separate paper (Nikutta et al., in preparation), we will extend our analysis of the cloud properties based on the X-ray data presented in this paper. ",
        "conclusions": "%  There has been much work so far into variable X-ray absorption in AGN. The accumulated research so far, including this work, has revealed eclipse events spanning a wide range of observed durations, ionization levels, and including full or partial covering.  Most of the short-term events (durations $\\la 1-3$ d) have been observed in the prominent case of the Sy 1.8 NGC~1365, where rapid and strong variations in $N_{\\rm H}$ are observed in a very large fraction of observations.  The clouds are inferred to exist in the BLR and are frequently modeled via variations in covering fraction of partial-covering material (e.g., De Rosa et al.\\ 2007; Turner et al.\\ 2008; Risaliti et al.\\ 2011).  In the context of clumpy-absorber models, there can exist a variable number of clouds in the line of sight at any given time, with a relatively extended X-ray continuum source located behind them.  \\textit{RXTE}, in contrast, has detected full-covering, cold or at most modestly ionized clouds.  An emerging picture is that clouds in the BLR and clouds in the torus may be two manifestations of the same radially extended cloud distribution, existing inside and outside, respectively, the dust sublimation region (Elitzur 2007).  That is, in the simplest picture, the full-covering clouds detected with \\textit{RXTE} are part of the same population of clouds as those causing the short-duration events, just occurring at larger radii, as we will demonstrate in $\\S$5.2.  It is potentially interesting that 4 of the 10 sources in our sample with secure or candidate eclipse events have multiple such events (certain objects thus seem to be prone to a higher frequency of events than others).  The high frequency of events observed in NGC~1365 may potentially be explained by a favorable geometry/viewing angle. However, NGC~1365 may be a special object; a statistical sample on a large number of AGN may be better suited to quantify eclipse events across all AGN.   \\subsection{Dependence of cloud events on AGN physical parameters}   %  Having identified secure and candidate eclipses in 10 of the 55 objects monitored with \\textit{RXTE}, we can attempt to explore if certain system parameters govern the presence or lack of eclipsing clouds.  For example, the 10 sources' black hole masses (see $\\S$5.2) span log($M_{\\rm BH}$) = 6.9--8.8, and these values are not extreme compared to black hole masses for Seyferts/quasi-stellar objects (QSOs) in general or our whole sample.  There is no evidence that having an extreme value of luminosity governs the presence/lack of eclipses: the range of log($L_{2-10}$) and log($L_{\\rm Bol}$) where eclipses are confirmed with \\textit{RXTE} to occur are 41.93--44.73 and 42.9--45.6, respectively.\\footnote{Values for $L_{\\rm Bol}$ for most sources are taken from Vasudevan et al.\\ (2010), except for Fairall~9 (Vasudevan \\& Fabian 2009), and MR~2251--178 and Cen~A ($L_{\\rm 2-10}$ from Rivers et al.\\ 2013 and $L_{\\rm Bol}/L_{2-10}$ corrections from Marconi et al.\\ 2004 for both objects).} Again, these values are not extreme.  In the context of clouds being produced in a disk wind, Elitzur \\& Ho (2009) estimate that such winds cannot exist when the bolometric luminosity drops below $5 \\times 10^{39}$ ($\\frac{M_{\\rm BH}}{10^7\\Msun}$)$^{2/3}$ erg s$^{-1}$. However, we cannot test this with our sample, as the lowest-luminosity objects monitored with \\textit{RXTE} have 2--10~keV luminosities of $2-6\\times 10^{40}$ erg s$^{-1}$ (NGC~4258, NGC~3998, and NGC~4051).  Accretion rate relative to Eddington, $L_{\\rm Bol}/L_{\\rm Edd}$, is also not an obvious factor; values of $\\dot{m}$ for the 10 objects span $\\sim0.1$ per cent for Cen~A up to $\\sim16$ per cent for MR~2251--178, with most other objects' values in the range of 3--10 per cent. These values are typical for Seyferts in general as well as for our sample.   Can radio loudness be a factor? One of our eight sources with secure eclipses is radio loud, compared to 13/55 (24 per cent) sources in the original sample. If we start with 55 objects, 13 of which are radio loud, and select at random eight of the 55 for a sub-sample, the probability $P(X=k)$ that exactly $k$ out of eight sources in the subsample will be radio loud is given by the hypergeometric distribution \\begin{equation} \\label{eq:hgd} P(X=k) = {K \\choose k} {N-K \\choose n-k} \\Big/ {N \\choose n}\\ , \\end{equation} where $N=55$ is the population size, $K=13$ is the number of successes in the population (radio-loud sources), and $n=8$ and $k$ are the sample size and the number of successes in the sample, respectively. ${\\cdot \\choose \\cdot}$ is a binomial coefficient. The chance of having exactly $k=0, 1, 2, 3$ or $4$ radio-loud sources in a random sample of eight is $P(X=0, 1, 2, 3,$ or $ 4) =$ (0.10, 0.29, 0.34, 0.20, or 0.07). In other words, it is almost as likely to have exactly 1/8 radio-loud sources in the sub-sample as it is to have 2/8 radio-loud sources. There is thus no significant statistical evidence that our sub-sample of sources with secure eclipses is unusual compared to the original sample in terms of fraction of radio-loud sources.  In summary, we find no evidence for a strong dependence of the presence of eclipse events as a function of the most common AGN parameters. One has to keep in mind that only strong correlations could have been detected with our relatively small number of observed events.  That being said, the reader is reminded that with a high dynamic range of sampling for 55 objects totaling 230 object-years of monitoring, this study is by far the largest available data set for testing the environment close to supermassive black holes via eclipse events over timescales longer than a few days.  \\subsection{Locations of X-ray obscuring clouds}    % \\label{sec:locations}  Constraints on the distance from the X-ray continuum source to each cloud $r_{\\rm cl}$ can be derived from the cloud's ionization level, column density, and eclipse duration.  Following, e.g., $\\S$4 of Lamer \\etal\\ (2003), we assume for simplicity that each cloud has a uniform density and ionization parameter, and that clouds are in Keplerian orbits. The cloud diameter $D_{\\rm cl}$ = $N_{\\rm}/n_{\\rm H} =  v_{\\rm cl}t_{\\rm D}$, where $t_{\\rm D}$ is the cloud's crossing time across the line of sight and $v_{\\rm cl}$ is the transverse velocity, equal to $\\sqrt{GM_{\\rm BH}/r_{\\rm cl}}$. Solving for $r_{\\rm cl}$ and using the definition of $\\xi$, one obtains (see, e.g., Eqn.\\ 3 of Lamer \\etal\\ 2003): $r_{\\rm cl} = 4\\times10^{16} M_7^{1/5} L_{42}^{2/5} t_{\\rm D}^{2/5} N_{\\rm H,22}^{-2/5} \\xi^{-2/5}$~cm, where $M_7 = M_{\\rm BH}/10^6 \\Msun$,   $L_{42}=L_{\\rm ion}/$($10^{42}$ erg s$^{-1}$), $N_{\\rm H,22}=N_{\\rm H}/$($10^{22}$~cm$^{-2}$), and $t_{\\rm D}$ is in units of days.  Table~\\ref{tab:ecl1dist} lists the inferred values of $r_{\\rm cl}$, obtained as follows: we use estimates of $M_{\\rm BH}$ from reverberation mapping where available; otherwise, estimates are from stellar kinematics or gas dynamics, based on empirical relations between optical luminosity and optical line widths, or from modeling reprocessing in accretion disks. The references are listed in column~4 of Table~\\ref{tab:ecl1dist}.  For all objects except Cen A, we estimate $L_{\\rm ion}$ as follows: we use the best-fitting unabsorbed power law from Rivers et al.\\ (2013), and extrapolate to the 0.1--13.6 keV range to estimate the luminosity $L_{\\rm 0.1-13.6 keV}$. The luminosity below $\\sim$0.1 keV, however, is expected to contain significant contributions from the thermal accretion disk continuum emission. Vasudevan \\& Fabian (2009) estimate that the ratio of the 0.0136--0.1 keV luminosity to the bolometric luminosity, $L_{\\rm 0.0136-0.1 keV}/L_{\\rm Bol}$, ranges from 0.2 to 0.6 for values of $L_{\\rm Bol}/L_{\\rm Edd}$ ranging from 0.01 to 0.6, respectively. We take values of $L_{\\rm Bol}$ and $L_{\\rm Bol}/L_{\\rm Edd}$ from Vasudevan et al.\\ (2010) when available. However, for Fairall~9 the value of $L_{\\rm Bol}$ is taken from Vasudevan \\& Fabian (2009), and for MR~2251--178, we use the 2--10 keV flux from Rivers et al.\\ (2013) and we estimate $L_{\\rm Bol}/L_{\\rm Edd} = 40$ from Marconi et al.\\ (2004). We then assume for simplicity a linear dependence of log10($L_{\\rm 0.0136-0.1 keV}/L_{\\rm Bol}$) on log10($L_{\\rm Bol}/L_{\\rm Edd}$) to estimate $L_{\\rm 0.0136-0.1 keV}$, and add that to $L_{\\rm 0.1-13.6 keV}$ to obtain the values of $L_{\\rm ion}$ listed in Table~\\ref{tab:ecl1dist}. The broadband SED of Cen~A is more like that of blazars than radio-quiet Seyferts, with a broad inverse-Compton hump dominating the emission from the optical band to gamma-rays. We assume $\\Gamma=1.84$ (Rivers et al.\\ 2013) over 0.4--13.6 keV, breaking to $\\Gamma=1.0$ over 0.0136--0.4 keV (Steinle 2010; Roustazadeh \\& B\\\"{o}ttcher 2011); this yields $L_{\\rm ion}$ = $10^{42.3}$ erg s$^{-1}$.  Our constraints on the ionization parameters for our clouds are poor. Given the column densities of our clouds and the energy resolution of the PCA, we can safely rule out values of log($\\xi$) above $\\sim2$. In that case, the absorber would manifest itself above 2~keV only via discrete lines and edges, as opposed to a broad rollover towards lower energies. We therefore calculate distances assuming log($\\xi$) = --1, 0, or +1.  Uncertainty on $r_{\\rm cl}$ is thus dominated by uncertainty on $\\xi$.  For brevity, we list in Table~\\ref{tab:ecl1dist} only values assuming log($\\xi$)=0.  Because $r_{\\rm cl} \\propto \\xi^{-2/5}$, values of $r_{\\rm cl}$ assuming e.g., log($\\xi$) = --1 or +1 are only factors of 2.5 larger or smaller, respectively. However, for Cen~A, we do not consider log($\\xi$) = +1. We fit the spectrum of the most absorbed spectral bin used in Rivers \\etal\\ (2011b), model the absorption with an \\textsc{xstar} table, and obtain log($\\xi$) $\\la -0.1$.  Similarly, for NGC~3227/2000--1, Lamer \\etal\\ (2003) found log($\\xi$) = --0.3--0 from a contemporaneous \\textit{XMM-Newton} observation.  These two events have among the highest quality energy spectra in the group, and one might speculate that such an ionization level may be representative of the rest of the sample.  For the candidate event in NGC~3516, we use the best-fitting ionization from Turner et al.\\ (2008), log($\\xi$)=2.19, as explained in Appendix A.  Finally, we note that although for a few events, it is difficult to estimate reliable uncertainties on $t_{\\rm d}$, a $\\pm10$ per cent uncertainty on $t_{\\rm d}$ translates into only $\\pm4$ per cent uncertainty in $r_{\\rm cl}$.   We obtain best-estimate values of $r_{\\rm cl}$ that are typically tens to hundreds of light-days from the central engine, and lie in the range $1 - 50 \\times 10^{4} R_{\\rm g}$ ($0.3 - 140 \\times 10^{4} R_{\\rm g}$ accounting for uncertainties), where $R_{\\rm g} \\equiv GM_{\\rm BH}c^{-1}$. We now compare these inferred radii to the locations of typical emitting components in AGN. In Table~\\ref{tab:BLRdust} we list inferred locations for the origins of optical and near-IR broad emission lines and ``hidden'' broad lines observed in polarized emission for type IIs, locations of IR-emitting dust determined via either interband correlations or IR interferometry, and locations of Fe K$\\alpha$ line-emitting gas determined via either X-ray spectroscopy or variability. For reverberation mapping of emission lines, we use time-lag results when available, otherwise we use full width at half-maximum (FWHM) velocities $v_{\\rm FWHM}$, assumed for simplicity Keplerian motion, and estimate radial locations via $r = GM_{\\rm BH}/( \\frac{3}{4} v_{\\rm FWHM}^2)$.  The results are also plotted in Fig.~\\ref{fig:plotrs} in units of $R_{\\rm g}$ for each object. All structures for a given object in Fig.~\\ref{fig:plotrs} are plotted on one dimension for clarity, but they do not necessarily overlap in space (e.g., X-ray clouds must lie along the line of sight, but other structures may lie out of it). The objects are listed in order of increasing $R_{\\rm d}$/$R_{\\rm g}$. The black points denote the estimates for $r_{\\rm cl}$ assuming log($\\xi$)=0, but we denote the minimum/maximum range assuming log($\\xi$)= --1 to +1 by the gray areas (again, with --1 to 0 for Cen~A and --0.3 to 0 for NGC~3227/2000--1).   The radial distance from the black hole where dust sublimates can be estimated via $R_{\\rm d} \\sim 0.4 (L_{\\rm bol}/10^{45} {\\rm erg s}^{-1})^{1/2} (T_{\\rm d}/T_{\\rm 1500})^{-2.6}~{\\rm pc}$, where $T_{\\rm d}$ is the dust sublimation radius, here assumed to be 1500 K (Eqn.\\ 2.1 of Nenkova et al.\\ 2008b). However, the boundary between dusty and dust-free zones is likely highly blurred, because relatively larger grains can survive at higher temperatures. In addition, individual components of dust can sublimate at different radii, e.g, graphite grains sublimate at a slightly higher temperature than silicate grains (Schartmann et al.\\ 2005). We use the above equation as an approximation for the outer boundary of the \"dust sublimation zone\" (DSZ), i.e., we assume that no dust sublimation occurs outside $R_{\\rm d}$. We take the inner boundary of the DSZ to be a factor of $\\sim$2--3 smaller than $R_{\\rm d}$. For example, Nenkova et al.\\ (2008b) point out that the $V$--$K$ band reverberation lags measured by Minezaki et al.\\ (2004) and Suganuma et al.\\ (2006) are $\\sim2-3$ times shorter than the light travel times predicted by the above equation. Those experiments thus may be tracking the innermost, larger grains. For simplicity, we assume that the central engine emits isotropically.\\footnote{This may not be true for Cen~A, whose continuum emission is likely mildly beamed. A typical value of Doppler $\\delta$ for Cen~A is 1.2 (Chiaberge et al.\\ 2001). Using 0.66 as the spectral index (Abdo et al.\\ 2010), the flux will be boosted by $\\delta^{2+\\alpha} \\sim 1.6$, which translates into only a $\\sim26$ per cent effect on our estimate of $R_{\\rm d}$.} The estimated values of $R_{\\rm d}$ are listed in Table~B1, with dusty regions represented by the purple areas in Fig.~\\ref{fig:plotrs}, and the inferred DSZs represented by the fading purple areas.  We list the radial locations of each cloud in units of $R_{\\rm d}$ in Table~\\ref{tab:ecl1dist}.  The clouds in 7/8 objects are consistent with residing in the DSZ considering uncertainties on $r_{\\rm cl}$. However, this ``clustering'' may be in part associated with our observational bias to select eclipses with events of $\\sim$tens of days, as per our selection function.  In Cen~A, the clouds are inferred to be consistent with residing entirely in the dusty zone. In contrast, NGC~5506 has the lowest value of $r_{\\rm cl}$/$R_{\\rm d}$; those clouds are likely the least dusty of the secure events in our sample.   The IR-emitting structures for Cen~A, NGC~3783, and NGC~3227 as mapped by either interferometric observations or optical-to-near-IR reverberation mapping are inferred to exist at radii $r_{\\rm IR} \\sim 3-20 \\times 10^4 R_{\\rm g}$. Although we can only make a statement based on three objects, values of $r_{\\rm cl}$ are generally consistent with these structures, again supporting the notion that the X-ray-absorbing clouds detected with \\textit{RXTE} are thus likely dusty.  For six of the seven objects with known BLRs (MR~2251--178, Mkn~509, Mkn~79, NGC~3783, NGC~3227, and Mkn~348), the BLR clouds are inferred to exist at $r_{\\rm BLR} \\sim 0.1 - 2 \\times 10^4 R_{\\rm g}$ from the black hole. We thus discuss our results in the context of the notion put forth by Netzer \\& Laor (1993) that the outer radius of the BLR may correspond to $R_{\\rm d}$ and the inner radius of the dusty torus, supported by the results of Suganuma et al.\\ (2006). In these six objects, one can see from Fig.~\\ref{fig:plotrs} that the radial ranges of $r_{\\rm BLR}$ are generally smaller than those for both $r_{\\rm cl}$ and $r_{\\rm IR}$. That is, best-estimate values of $r_{\\rm cl}$ are generally commensurate with the outer portions of the BLR or exist at radii up to $\\sim$15 times that of the outer boundary of the BLR, although we must caution that we are dealing with a very small sample of only six objects.  We can conclude that, at least for these six objects, the X-ray-absorbing clouds are likely more consistent with the dusty torus than with the BLR. Note that this conclusion is relatively robust to our assumed values of log($\\xi$); even if log($\\xi$)=+2, the X-ray-absorbing clouds would be closer to the black hole by only a factor of 2.5. Our results thus demonstrate that radii commensurate with the outer BLR and with the inner dusty torus are radii where X-ray-absorbing clouds may exist.   The seventh object with measured BLR parameters is NGC~5506, where $r_{\\rm BLR} \\sim 4-10 \\times 10^4 R_{\\rm g}$ based on IR emission lines. The inferred radial location of the X-ray cloud is $0.3-4.5\\times 10^4 R_{\\rm g}$ -- roughly commensurate with the innermost BLR or distances slightly smaller than the BLR as traced by IR emission lines, and thus likely non-dusty. The eclipse event in NGC~5506 may thus be analogous to short-term ($t_{\\rm D} \\la$ 1 d) eclipse events observed in other objects, namely NGC~1365 and Mkn~766 (Risaliti et al.\\ 2005, 2007, 2009b, 2011; Maiolino et al.\\ 2010). The variable absorbers detected in those observations have column densities typically $N_{\\rm H} \\sim 1-5 \\times 10^{23}$~cm$^{-2}$ and can be partial covering, as for Mkn~766. They are inferred to exist $\\sim10^{2.5} - 10^4 R_{\\rm g}$ from the black hole, and are frequently identified as BLR clouds.    We return to a point regarding the type Is from $\\S$4.3, regarding eclipse events longer than $\\sim100$~d in duration.  Our selection function analysis shows that if such events occurred in type Is with the same frequency as in type IIs, we should have detected them.  It is thus possible that type Is are devoid of clouds at these radial distances along the line of sight.  Another possibility is that these clouds do exist at these radii in type Is, but intersect our line of sight so rarely [probabilities $<$ 1/(10--300)] that 16 years of monitoring 37 type I AGN was insufficient to catch even one event (excluding the possibility that the event in MR~2251--178 falls into this category).  In either case, we can estimate the corresponding radial distances.  We assume a black hole mass of 10$^{7.7} \\Msun$, $N_{\\rm H} = 12 \\times 10^{22}$~cm$^{-2}$, log($\\xi$)=0, and log($L_{\\rm ion}$) = 44.3, the non-weighted average of the type Is with secure events.  Arbitrarily-chosen durations of 100, 316, 1000, and 3162 d correspond to inferred radial distances of 14, 23, 36, and 57 $\\times 10^4$ $R_{\\rm g}$.  Finally, we investigate any possible trend between the best-fitting values of peak $N_{\\rm H}$ and the inferred values of $r_{\\rm cl}$ listed in Table~\\ref{tab:ecl1dist}, in units of either light-days, $R_{\\rm g}$, or $r_{\\rm cl}$/$R_{\\rm d}$ (Fig.~\\ref{fig:NHvsrad}). Given the small number of eclipse events, and given that values of $N_{\\rm H}$ span only a factor of $\\sim$7, we find no robust correlations; inclusion/exclusion of only 1--2 secure or candidate points significantly changes the resulting % coefficients and null hypothesis probabilities.  Note that the overdensity of eclipse events around 200 light-days (left-hand panel), $8\\times10^4 R_{\\rm g}$ (middle panel), and $\\la 1 R_{\\rm d}$ (right-hand panel) is likely not intrinsic to the sources but mainly caused by our selection function. This figure should thus not be used to derive constraints on the radial distribution of clouds.  \\begin{table*} \\begin{minipage}{170mm} \\caption{Inferred cloud distances from the black hole}   % \\label{tab:ecl1dist} \\begin{tabular}{llcccccccc} \\hline Source   &           & log($M_{\\rm BH}$) &       & $L_{2-10}$, $L_{\\rm ion}$, $L_{\\rm Bol}$, & log($\\xi$, erg         &  $r_{\\rm cl}$      &  $r_{\\rm cl}$    &                       &  \\\\ name     & Event     &  ($\\Msun$)      & Ref.\\ & (erg s$^{-1}$)          & cm s$^{-1}$)  & (light-days)          & ($10^4$ $R_{\\rm g}$) & $r_{\\rm cl}$/$R_{\\rm d}$ &  Notes \\\\ \\hline NGC 3783 & 2008.3        & 7.47        & VP06 &  43.2, 44.1, 44.4       &  0                & $147^{+11}_{-10}$  & $8.6^{+0.7}_{-0.6}$   & $0.62\\pm0.04$         &     \\\\ & (Full event)  &             &      &                         &                   &                  &                    &                        &   \\medskip    \\\\ Mkn 79   & 2003.5        &  7.72       & Pe04 &  43.3, 44.3, 44.7       &  0                & $229^{+83}_{-79}$  & $7.5\\pm2.7$        & $0.68\\pm0.24$           &   \\\\ & 2003.6        &             &      &                         &  0                & $290^{+42}_{-35}$  & $9.6^{+1.4}_{-1.1}$   & $0.86^{+0.13}_{-0.09}$     &   \\\\ & 2009.9        &             &      &                         &  0                & $314^{+91}_{-74}$  & $10.4^{+3.0}_{-2.5}$   & $0.93^{+0.27}_{-0.22}$     &   \\medskip \\\\ Mkn 509  & 2005.9        &  8.19       & VP06 &  44.3, 44.9, 45.2       &  0                & $851^{+255}_{-278}$ & $9.5^{+2.9}_{-3.1}$  & $1.42^{+0.42}_{-0.46}$     &  $^{\\dagger}$ \\medskip \\\\ MR 2251--178& 1996  & 8.3 & W09  &  44.7, 45.3, 45.6          &  0                & 460--5700        & 4--49                   &  0.5--6.0              &   $^{\\dagger}$$^{\\ddagger}$ \\medskip \\\\ NGC 3227 & 2000--1       &  6.88       & D10  &  42.5, 43.0, 43.5       &  --0.3            & $82^{+9}_{-8}$     & $19\\pm2$           &  $0.97\\pm0.11$         &    \\\\ &               &             &      &                         &  0                & $62^{+7}_{-6}$     & $14\\pm1$           &  $0.74\\pm0.05$         &  \\\\ % & 2002.8        &             &      &                         &  0                & $23\\pm7$         & $5.3\\pm1.6$        &  $0.28^{+0.08}_{-0.09}$   &  \\medskip \\\\ Cen~A    & $\\sim$2003--4 &  7.78       & R11  &  41.9, 42.3, 42.9       &  0                & $214^{+70}_{-93}$   & $6.2^{+2.0}_{-2.6}$  & $5.3^{+1.6}_{-2.1}$      &   $^{\\dagger}$ \\\\ & 2010--1       &             &      &                         &  0                & $101^{+7}_{-6}$    & $2.9\\pm0.2$        & $2.4\\pm0.2$           &   \\medskip \\\\ NGC 5506 & 2000.2        &  7.94       & Pa04 &  43.0, 43.9, 44.3       &  0                & $62^{+26}_{-24}$    & $1.2^{+0.6}_{-0.4}$  & $0.26^{+0.11}_{-0.08}$    &  \\medskip \\\\ Mkn 348  & 1996--7       &  7.18       & WZ07 &  43.0, 44.1, 44.4       &  0                & $432^{+79}_{-73}$    & $50\\pm9$          & $1.82\\pm0.37$       &   $^{\\dagger}$  \\medskip \\\\ Fairall 9 & 2001.3 (cand.) & 8.41      & VP06 &  43.4, 44.6, 44.8       &  0                & $>$180            & $>$1.2            & $>0.5$              &   $^{*}$ \\medskip  \\\\ NGC 3783 & 2008.7 (cand.)& 7.47        & VP06 &  43.2, 44.1, 44.4       &  0                & $>$230            & $>$13.7            & $>1.0$              &   $^{*}$   \\\\ & 2011.2 (cand.)&             &      &                         &  0                & $127^{+71}_{-52}$   & $7.4^{+4.2}_{-3.0}$   & $0.54^{+0.30}_{-0.22}$    &   \\medskip   \\\\ NGC 3516 & 2011.7 (cand.) & 7.50       & D10  & 43.0, 43.7, 44.2        & +2.19             & 34                & 1.9               & $\\sim0.6$           & $^{**}$ \\\\      \\hline  \\multicolumn{2}{c}{Average for 12 Secure Events} & 7.63  &    &    43.1, 43.8, 44.2 &               & 200             & 8.0               & 1.0                 &   \\\\ \\hline \\end{tabular}\\\\ Estimates of the distance from the central black hole to each eclipsing cloud $r_{\\rm cl}$ following $\\S$5.2.  Uncertainties listed above do not include uncertainty on the ionization parameter $\\xi$. Reasonable ranges for log($\\xi$) are --1 to +1 for most objects, which translates into factors of 2.5 larger/smaller (for Cen~A, values of log($\\xi$) ranging from --1 to 0 are plausible; for NGC~3227/2000--1, Lamer et al.\\ (2003) measured log($\\xi$) = --0.3--0). $L_{2-10}$ is the 2--10~keV of the hard X-ray power-law component from the long-term time-averaged \\textit{RXTE} spectrum (Rivers \\etal\\ 2013). $L_{\\rm ion}$ is the 1--1000 Ryd ionizing luminosity; please see $\\S5.2$ for details. $L_{\\rm Bol}$ is the bolometric luminosity; please see $\\S5.1$ for details. Luminosities are corrected for all (intrinsic and Galactic) absorption. $R_{\\rm d}$ denotes the outer boundary of the ``dust sublimation zone'' (Nenkova et al.\\ 2008b), i.e., dust residing at distances greater than $R_{\\rm d}$ likely does not sublimate, while distances smaller than $\\sim\\frac{1}{2}-\\frac{1}{3}R_{\\rm d}$ are expected to be dust-free. Average values (last row) were determined in log space. References for $\\Msun$ are: D10 = Denney \\etal\\ (2010), Pa04 = Papadakis (2004), Pe04 = Peterson \\etal\\ (2004), R11 = References in Rothschild \\etal\\ (2011), VP06 = Vestergaard \\& Peterson (2006), W09 = Wang \\etal\\ (2009), and WZ07 = Wang \\& Zhang (2007).   \\\\ $^{\\dagger}$Estimates of $r$ may be lower if peak $N_{\\rm H}$ was higher than that measured because either the beginning or end of the eclipse was not covered by monitoring data. \\\\ $^{\\ddagger}$For MR~2251--178, we have only limits to the duration, but no reliable ``best estimate.''\\\\ $^{*}$Upper limits on $N_{\\rm H}$ were used to derive lower limits on $r_{\\rm cl}$. \\\\ $^{**}$Assumed ionization following Turner et al.\\ (2008); see Appendix A. \\end{minipage} \\end{table*}    \\begin{figure*} \\includegraphics[angle=-90,width=1.00\\textwidth]{fig8.eps}   % \\caption{Plot of inferred radial locations of the secure X-ray-absorbing clouds $r_{\\rm cl}$, plus inferred radial locations of IR continuum-emitting dust and BLR clouds, in units of $R_{\\rm g}$. See Table~\\ref{tab:BLRdust} for details.  Sources are ranked in order of increasing $R_{\\rm d}$/$R_{\\rm g}$.  Locations of each of the X-ray absorbing clouds for an assumed ionization value of log($\\xi$)=0 are denoted by black points, with the full range obtained with log($\\xi$) spanning values from --1 to +1 (with --1 to 0 for Cen~A and --0.3 to 0 for NGC~3227/2000--1) indicated by the gray shaded areas. The data point for MR~2251--178 is omitted due to the poor constraints on $t_{\\rm D}$ and thus $r_{\\rm cl}$.  The purple shaded areas denote radial distances expected to contain dust, with the fading purple areas (1.0 down to 0.4$R_{\\rm d}$) denoting the ``dust sublimation zone.''  Known locations of IR continuum emitting dust are plotted in red points, with ranges indicated by pink shaded areas.  BLR ranges are indicated by the yellow areas.  Green points indicate origins of UV broad emission lines (e.g., He~\\textsc{ii}, C~\\textsc{iv}, S~\\textsc{iv}). Orange points show origins of optical broad emission lines (e.g., Balmer series). Magenta points show origins of IR broad emission lines (e.g., Bracken and Paschen series, He~\\textsc{i}$\\lambda$2.058$\\mu$). Locations of Fe K$\\alpha$-emitting gas are plotted in cyan.  As a reminder, radial locations of these structures may be commensurate but structures may not physically overlap, e.g., the X-ray-absorbing clouds lie on the direct line of sight to the central engine, but emission components may lie off the line of sight.} \\label{fig:plotrs} \\end{figure*}   \\begin{figure*} \\includegraphics[angle=-90,width=1.00\\textwidth]{fig9.eps} % \\caption{Best-fitting values of $N_{\\rm H}$ plotted against inferred radial location $r_{\\rm cl}$ in units of light-days (left), in units of gravitational radii $R_{\\rm g}$ (center; $R_{\\rm g} \\equiv GM_{\\rm BH}/c^2$), and relative to dust sublimation radius $R_{\\rm d}$. Filled and open symbols correspond to type I and II objects, respectively. Colors are the same as in Fig.~\\ref{fig:durbars}. } \\label{fig:NHvsrad} \\end{figure*}    \\subsection{A possible double-eclipse event in NGC~3783}   % \\label{sec:3783double} The $N_{\\rm H}$($t$) and $HR1$ profiles for the NGC~3783/2008.3 event are highly intriguing, as the two spikes suggest two absorption events separated by only 11~d.  The fact that, assuming that $\\Gamma$ remains constant during this time, $N_{\\rm H}$($t$) remains near $4\\times 10^{22}$~cm$^{-2}$ in between the spikes is also intriguing. No such complex $N_{\\rm H}$($t$) profile has been observed for any other AGN to date.  As illustrated in Fig.~\\ref{fig:3783doubleprofile}, we model the $N_{\\rm H}$($t$) profile in several ways, fitting only bins \\#5--14 (see Table~\\ref{tab:3783TRtable} for bin definition; refers to those bins with non-zero values of $N_{\\rm H}$ plus two bins with upper limits before and after).  We first test a single uniform-density sphere (gray line), which yields $\\chi^2/dof$ = 40.1/11=3.65 and underestimates the spikes in $N_{\\rm H}$ in bins \\#9 and \\#15.  We also fit the profile assuming two independent eclipsing events.  We test a profile from two uniform-density spheres (cyan line), and one from two linear-density spheres profile (red line). For the latter, the density profile is $n$($r$) = $n_{\\rm max}$(1--$r/R$), where $n_{\\rm max}$ is the number density at the center and $R$ is the outer radius of the cloud along the transverse axis (red line).  Each cloud's center point in time space is held frozen at the midpoints of bins \\#9 or 15.  The best-fitting models yield $\\chi^2/dof$ = 50.2/10=5.02 and 48.9/10=4.89, respectively, and did not correctly fit bin \\#9 nor the $\\sim$2--4 bins between the spikes.  We also fit each spike with a phenomenological $\\beta$-profile of the form $N_{\\rm H}$($t$) = $N_{\\rm H, max} \\sqrt{1 - r/R_{\\rm c}}$ (panel c), where $N_{\\rm H, max}$ is the column density along a line of sight through the cloud's center and $R_{\\rm c}$ is the core radius (Dapp \\& Basu 2009).  The best fit yields $\\chi^2/dof$ = 39.5/10=3.95 but still underestimates the middle part of the profile.   Furthermore, the probability that two separate eclipse events could independently occur so close together in time is low: NGC~3783 was monitored for 3.27~yr, had no observations for 1.88~yr, and then was monitored for an additional 7.84~yr.  We perform Monte Carlo simulations wherein we assume a 1.0-d grid and we randomly place three eclipses in the 3.27 or 7.84-yr campaigns, and empirically estimate the probability that any two of their peaks can occur $<$15 d apart: only 0.24 per cent.  Combined with the fact that $N_{\\rm H}$($t$) does not return to zero between the two spikes, the likelihood that the profile is comprised of two independent eclipse events seems low. The addition of a \\textit{third} independent cloud with $N_{\\rm H} \\sim 4 \\times 10^{22}$~cm$^{-2}$ transiting the line of sight roughly halfway between the maximum transits of the other two would of course yield a good fit to the observed $N_{\\rm H}$($t$) profile, but this is statistically highly unlikely ($\\sim5\\times 10^{-6}$).   Finally, we model a uniform-density shell (green line in Fig.~\\ref{fig:3783doubleprofile}; subtracting a smaller uniform-density sphere from another) to successfully model fit the middle part and the two spikes.  The best fit has $\\chi^2/dof$ = 13.2/9=1.47, with the outer and inner shell boundaries taking $7.5^{+0.5}_{-1.0}$ and $4.5\\pm0.5$~d, respectively, to transit the line of sight.  However, we stress that this fit is purely phenomenological; it is not clear how such a structure could be created or survive for long periods in an AGN environment.   We speculate that the observed profile could be caused by two dense clumps with some sort of connecting structure (a dumbbell shape), with parts aligned along the direction of travel across the line of sight. We discuss a possible physical origin for this behavior in $\\S$\\ref{sec:profiles}.   \\begin{figure} \\includegraphics[angle=-90,width=0.48\\textwidth]{fig10.eps}      % \\caption{$HR1$ and $N_{\\rm H}$($t$) light curves for the NGC~3783/2008.3 event, with its suspected ``dual-clump'' profile. The solid lines indicate the best-fitting models to the $N_{\\rm H}$($t$) profile; we fit only bins \\#5--14 (those bins with non-zero values of $N_{\\rm H}$ plus two bins with upper limits before and after; see Table~\\ref{tab:3783TRtable} for bin definition). We test a single uniform-density sphere (gray), a dual uniform-density sphere model (cyan), a dual linear-density-sphere model, a dual $\\beta$-profile (orange), and a uniform-density shell (green).  Values of $\\chi^2/dof$ are in parentheses.} \\label{fig:3783doubleprofile} \\end{figure}   \\subsection{Physical properties of obscuring clouds}     % \\label{sec:diams}    We estimate diameters, number densities and masses for each eclipsing cloud using the equations in $\\S$\\ref{sec:locations} and listed them in Table~\\ref{tab:clumpsizemass}. The average diameter across the sample of secure events is $\\sim0.25$ light-days, with an average number density of $\\sim3\\times10^8$~cm$^{-3}$. Nenkova \\etal\\ (2008b) estimate theoretically that a total of $\\sim10^4 - 10^5$ clouds comprises the dusty torus. The mass estimates for individual clouds in Table~\\ref{tab:clumpsizemass} span $\\sim10^{-8}-10^{-2} \\Msun$ with a mean value in log space of $10^{-5.2} \\Msun$, % thus suggesting a total mass for the clumpy component of the torus (i.e., excluding any thin accretion disk or intercloud diffuse medium) of $\\sim 10^{-4} - 10^{3} \\Msun$.  % If the torus clouds are produced in a disk wind, then this estimate indicates the maximum mass associated with the outflow.  Since the value of $\\sim10^4 - 10^5$ clouds quoted above likely refers only to IR-emitting clouds located outside $R_{\\rm d}$, this mass refers to clouds lying outside $r_{\\rm d}$ only, and also ignores contributions from non-clumpy material such as any intercloud medium (Stalevski et al.\\ 2012).    \\begin{table*} \\begin{minipage}{145mm} \\caption{ Inferred physical properties of each cloud; inferred size of X-ray-emitting region }   % \\label{tab:clumpsizemass} \\begin{tabular}{llccccc}  \\hline Source    &              &  Diam.            & Diam.                          & Num.\\ dens., $n$   & log(mass)             & $D_{\\rm X-src}$   \\\\ name      & Event        & ($10^{14}$~cm)     & (light-days)                       & ($10^8$~cm$^{-3}$)  & ($\\Msun$)         & ($R_{\\rm g}$)       \\\\   \\hline NGC~3783  & 2008.3       & $1.3^{+0.7}_{-0.5}$  & $0.048^{+0.029}_{-0.018}$         & $8.6^{+7.0}_{-3.9}$   & $-6.0^{+0.4}_{-0.5}$     &  $<47$      \\\\ Mkn~79    & 2003.5       & $2.4^{+2.6}_{-1.5}$  & $0.089^{+0.097}_{-0.057}$          & $6.0^{+15.9}_{-4.0}$  & $-5.4^{+0.5}_{-0.8}$    &  $<64$  \\\\ Mkn~79    & 2003.6       & $3.1^{+1.8}_{-1.2}$  & $0.11^{+0.07}_{-0.04}$            & $3.7^{+3.9}_{-1.9}$   & $-5.3\\pm0.5$          & $<65$    \\\\ Mkn~79    & 2009.9       & $2.4^{+2.1}_{-1.3}$  & $0.088^{+0.078}_{-0.046}$            & $3.2^{+5.5}_{-2.0}$   & $-5.7^{+0.4}_{-0.5}$     & $<57$    \\\\ Mkn~509   & 2005.9       & $4.9^{+5.7}_{-3.3}$  & $0.18^{+0.21}_{-0.12}$            & $1.8^{+4.5}_{-1.1}$   & $-5.1^{+0.6}_{-0.9}$     & $<46$    \\\\ MR~2251--178 & 1996       & $0.25-91$          &  $0.009-3.4$                  & $0.06-30$           & $-7.7$ to $-2.7$     & $<310$    \\\\ NGC~3227  & 2000--1       & $5.8^{+0.3}_{-1.0}$   &  $0.22^{+0.01}_{-0.04}$          & $3.9^{+1.4}_{-0.8}$    &  $-4.5^{+0.1}_{-0.2}$    & $<540$    \\\\ NGC~3227  & 2002.8        & $0.49^{+0.56}_{-0.31}$ & $0.018^{+0.020}_{-0.012}$         & $27^{+62}_{-17}$      &  $-6.9^{+0.6}_{-0.8}$    & $<92$    \\\\ Cen A     & $\\sim$2003--4 & $120^{+60}_{-90}$     & $4.5^{+2.3}_{-3.4}$              & $0.07^{+0.23}_{-0.03}$ & $-2.3^{+0.3}_{-0.2}$    & $<2000$    \\\\ Cen A     & 2010--1       & $26.9^{+0.1}_{-10.1}$  & $1.00^{+0.01}_{-0.38}$           & $0.30^{+0.24}_{-0.04}$   & $-3.6^{+0.1}_{-0.5}$    & $<300$    \\\\ NGC 5506  & 2000.2        & $0.12^{+0.13}_{-0.08}$ & $4.3^{+5.8}_{-2.9}\\times10^{-3}$ & $34^{+111}_{-23}$     & $-8.6^{+0.5}_{-0.9}$    & $<1.9$  \\\\ Mkn~348   & 1996--7       & $20^{+17}_{-10}$      & $0.74^{+0.62}_{-0.36}$           & $0.90^{+1.17}_{-0.49}$  & $-3.5\\pm0.5$        &  $<1600$ \\\\ Fairall~9 & 2001.3        & $<4.5$              & $<0.17$                       & $\\sim 5-27$       & $\\sim-6.2$ to  $-4.7$  & $<12$\\\\ NGC~3783  & 2008.7        & $<2.8$              & $<0.10$                       & $\\sim 1-5$        & $\\sim-7.0$ to $-5.9$   & $<63$ \\\\ NGC~3783  & 2011.2       & $0.9^{+1.0}_{-0.6}$  & $0.035^{+0.035}_{-0.023}$         & $12^{+41}_{-8}$       & $-6.4^{+0.4}_{-0.8}$      & $<43$   \\\\ NGC~3516  & 2011.7        & $\\sim11$            & $\\sim0.4$                     & $\\sim 0.4$        & $\\sim-4.6$              & $<230$\\\\  \\hline \\multicolumn{2}{c}{Average for 12 Secure Events} &   3.9  & 0.25                & 2.6               & --5.2                    &  $<120$  \\\\ \\hline  \\end{tabular}\\\\ Best-estimate values listed here correspond to best-fitting values of $N_{\\rm H}$ and values of duration listed in Table~\\ref{tab:ecl1summ} and an assumed ionization state of log($\\xi$)=0. Uncertainties listed here are based on the uncertainties on observed duration, $N_{\\rm H}$, and the ranges of $r_{\\rm cl}$ and log($\\xi$) listed in Table~\\ref{tab:ecl1dist}. $D_{\\rm X-src}$ denotes inferred upper limits on the size of the X-ray continuum source, assuming that the cloud completely covers the X-ray continuum source. Average values were determined in log space. \\end{minipage} \\end{table*}   Finally, assuming that each eclipsing cloud fully covers the X-ray continuum source behind it, the upper limits on the diameter of the cloud yield inferred upper limits on the size of the X-ray continuum source, $D_{\\rm X-src}$. These limits are listed in Table~\\ref{tab:clumpsizemass} in units of $R_{\\rm g}$. Among the secure events, the limits are as low as a few tens of $R_{\\rm g}$ for NGC~3783, Mkn~79, and Mkn~509, and even as low as $2R_{\\rm g}$ in the case of NGC~5506.  The X-ray continuum in Seyferts is generally thought to originate via inverse Compton scattering of the soft UV photons from the accretion disk by hot electrons in a compact corona, possibly located in the innermost accretion disk or at the base of a jet (e.g., Haardt \\etal\\ 1997; Markoff \\etal\\ 2005). X-ray microlensing analysis of distant quasars typically yields upper limits on the size of the X-ray corona of $\\sim10R_{\\rm g}$ (e.g., Chartas \\etal\\ 2012, Chen \\etal\\ 2012, and references therein). Recent results using X-ray reverberation lags imply sizes of a few to $\\sim10$ $R_{\\rm g}$ (Reis \\& Miller 2013). Geometrical limits from absorption variability in NGC~1365 typically yield $D_{\\rm X-src} \\la 30-60 R_{\\rm g}$ (Risaliti \\etal\\ 2007, 2009b). Our estimates of $D_{\\rm X-src}$ are consistent with these studies.   \\subsubsection{Exploring the range in observed density profiles} \\label{sec:profiles}  We speculated in $\\S$\\ref{sec:3783double} that the unusual density profile of the NGC~3783/2008.3 event could be caused by the cloud having a dumbbell shape. Even more speculatively, perhaps this structure is a cloud that is in the process of getting torn in half. As pointed out by Krolik \\& Begelman (1988), the self-gravity of a single cloud in the vicinity of a supermassive black hole is not highly effective against tidal shearing, and clouds can get significantly stretched out within an orbit. Resistance to tidal shearing requires that the size of the cloud is limited to $\\la D_{\\rm shear} = 10^{16} N_{\\rm H,23} r_{\\rm pc}^3 M_7^{-1}$~cm, where $N_{\\rm H,23} = N_{\\rm H}/10^{23} {\\rm cm}$, $r_{\\rm pc}$ is the distance from the black hole to the cloud, and $M_7 = M_{\\rm BH}/(10^7 \\Msun)$ (Elitzur \\& Shlosman 2006). For NGC~3783/2008.3, $D_{\\rm shear} \\sim 7\\times10^{12}$~cm, a factor of $\\sim1/18$ times the inferred cloud diameter $D_{\\rm cl}$, so the possibility of tidal shearing seems plausible. In fact, for all secure eclipse events, we obtain values of $D_{\\rm cl}/D_{\\rm shear}$ spanning 2.3--350, with an average value (in log-space) of 26, suggesting that many of these clouds' sizes put them at risk of being tidally disrupted or sheared.  The unusual density profile could also be consistent with models wherein torus clouds may be in the form of clumpy winds, as opposed to compact, discrete clouds, originating at the accretion disk.  These winds may be magnetohydrodynamic (MHD) driven (e.g., Konigl \\& Kartje 1994; Fukumura et al.\\ 2010), with local column densities that can be consistent with values measured in this paper.  Other models feature IR radiation-driven winds (Dorodnitsyn \\& Kallman 2012; Dorodnitsyn et al.\\ 2012): UV/X-ray photons are reprocessed into IR thermal emission in the torus, and IR radiation pressure vertically drives dust grains. Moderately Compton-thick flows are possible at distances of $\\ga$1 pc from the black hole (if our assumption about Keplerian motion is not valid, then distance estimates for our clouds may be in error).  The models generally assume that clouds' self-gravity is negligible.  Similarly, Czerny \\& Hryniewicz (2011) proposed a turbulent, dusty, disk wind as the origin for the low-ionization region of the BLR: clouds rise from the disk, intense external radiation heats the gas and dust sublimates.  In all these models, the torus is highly dynamic and can feature cloud motions perpendicular to the disk (including failed winds), and thus not strictly Keplerian.  In any case, the $N_{\\rm H}$($t$) profile of the NGC~3783/2008.3 event can thus also be explained via a non-homogeneous mass outflow, featuring two overdensities that crossed the line of sight 11~d apart.  However, the profile of NGC~3783/2008.3 contrasts with three of the other eclipse events that have $N_{\\rm H}$($t$) profiles that are also well resolved in time but are symmetric and centrally peaked: NGC~3227/2000--1 and Cen~A/2010--1, observed with \\textit{RXTE}, and SWIFT J2127.4+5654 (Sanfrutos \\etal\\ 2013). Such profiles pose a challenge to MHD/IR-driven wind models. For example, the IR-driven winds of Dorodnitsyn et al.\\ (2012) suggest that $\\Delta$$N_{\\rm H}$/$N_{\\rm H}$ is typically $\\sim$ a few (less than the factors we observe). In addition, such winds are inferred to originate $\\sim1-3$ pc from the black hole, farther than the distances inferred for our observed eclipses.  In summary, the self-gravity of these clouds likely dominates over any tidal shearing/internal turbulence, presenting a challenge to the expectation that clouds are easily tidally sheared. In these cases, one or more of the assumptions that go into the calculations of $D_{\\rm cl}$ or $D_{\\rm shear}$ may be wrong.  Alternately, some other physical process may prevent clouds from shearing, e.g., confinement of these clouds by external forces (e.g., gas pressure from the ambient medium, external magnetic fields; Krolik \\& Begelman 1988) may be important.   \\subsection{Implications for clumpy-absorber models}  %  \\C\\ torus models have found a lot of observational support, particularly from recent IR observations. For example, smooth-density torus models always predict that the silicate features at 9.7 and 18\\mic\\ will be in emission for pole-on viewing and in absorption for edge-on viewing. However, samples of mid-IR spectra of type I Seyferts can exhibit silicate features spanning a range of emission \\textit{and} absorption (see Hao et al.\\ 2007 for a sample obtained with \\textit{Spitzer}); type IIs generally show only weak absorption but a few can show emission (Sturm et al.\\ 2005; Hao et al.\\ 2007). Nenkova et al.\\ (2008b) and Nikutta et al.\\ (2009) demonstrate that the clumpy-torus models are consistent with these observations and can explain mismatches between the optical classification and the expected behavior of the silicate features. Clumpy-torus models predict nearly isotropic mid-IR continuum emission and anisotropic obscuration, as generally observed (e.g., Lutz et al.\\ 2004; Horst et al.\\ 2006). Finally, clumpy tori have been predicted to host both hot and cooler dust in close proximity to each other (Krolik \\& Begelman 1988), in contrast to smooth-density tori. Recent IR interferometric observations indeed do affirm such a co-existence of hot (800~K) and cooler ($\\sim 200-300$~K) dust components in nearby AGN (Jaffe et al.\\ 2004; Poncelet et al.\\ 2006; Tristram et al.\\ 2007; Raban et al.\\ 2009).   In this subsection, we derive the first \\textit{X-ray} constraints on the \\C\\ model parameters, thus obtaining constraints in a manner independent from IR SED fitting.  We caution, however, that there may not be an exact correspondence between constraints derived from the two methods, since IR emission arises only from dusty clouds outside the DSZ while our X-ray clouds can be dust free.  The clumpy-torus models are generally defined using the following free parameters: \\tv\\ is the $V$-band optical depth of single clouds; usually, all clouds are assumed to have identical values of \\tv. The radial extent of the dusty torus is characterized by $Y$, the ratio of the outer radial extent of clouds to the dust sublimation radius $R_{\\rm d}$. $\\Ni(\\theta)$ is the average number of clouds along a radial line of sight that is an angle $\\theta$ from the equatorial plane. Assuming a Gaussian distribution, $\\Ni(\\theta,\\sigma,\\No) = \\No \\exp(-(\\theta/\\sigma)^2)$, where \\No\\ is the average number of clouds along a radial line of sight in the equatorial plane (typically 5--15), and $\\sigma$ parametrizes the angular width of the cloud distribution. The index $q$ describes the radial dependence of the average number of clouds per unit length, $\\Nc(r,q) \\propto r^{-q}$.  The inclination angle of the system is $i \\equiv 90\\degr-\\theta$, defined such that $i=0$ ($\\theta=90\\degr$) denotes a system with its equatorial plane in the plane of the sky.  The predicted shape of the IR SED is sensitive to all of these parameters. Ramos Almeida et al.\\ (2009, 2011; hereafter RA09 and RA11) have successfully fit the mid-IR spectra of $\\la20$ AGN with \\C\\ torus models and the \\textsc{BayesClumpy} tool (Asensio Ramos \\& Ramos Almeida 2009). There are only three sources in our eclipse sample that overlap with the samples of RA09 and RA11 (NGC~3227, Cen~A, and NGC~5506). In the following we attempt to focus on comparing our derived parameter constraints for the type I/II classes as opposed to for individual objects.   $\\bullet$ \\tv: we can translate our inferred values of $N_{\\rm H}$ for individual clouds into values of $\\tv$, which we can then compare to those values typically used in the $\\C$ model fitting.  From the ``secure'' clumps we observe, $\\Delta$$N_{\\rm H} = 4-26 \\times 10^{22}~{\\rm cm}^{-2}$, with values spanning very similar ranges for both type Is/IIs. Using the Galactic dust/gas conversion ratio, $N_{\\rm H} = A_{\\rm V} {\\rm (mag)}  \\times 1.8 \\times 10^{21} {\\rm cm}^{-2}$ (Predehl \\& Schmidt 1995), these values translate to $A_{\\rm V} = 22-144$ mag, or $\\tv = 20-132$.  The \\C\\ models use one value of $\\tau_{\\rm V}$ per object/spectrum to represent the weighted average value across all clouds in that object.  In contrast, we probe 1--3 individual clouds per object.  Subject to small-number statistics, we can therefore test both the inter-object scatter of \\tv\\ values, as well as the scatter from cloud to cloud in one given source. From general theoretical considerations, Nenkova et al.\\ (2008b) suggest that \\tv\\ be typically 30--100. RA11, who use the entire \\C\\ database of models in their IR SED fitting, estimate values of 40--140 for type Is and 5--95 for type IIs. Alonso-Herrero et al.\\ (2011), with the same approach but different sources, find \\tv\\ = 105--146 (type Is) and 49--130 (type IIs). These IR-derived ranges are fully consistent with our X-ray based values. In contrast to the above studies, we find no evidence for different cloud optical depths between type I and II sources.  Three of our sources are also present in the IR samples of RA11 and Alonso-Herrero et al.\\ (2011). In two of these (Cen~A and NGC~3227) our X-ray constraints agree very well with the IR ones. However, our value for NGC~5506 is smaller than that in RA11 by a factor of $\\sim2$ and smaller than that of Alonso-Herrero et al.\\ (2011) by $\\sim5$. For our sources with multiple eclipsing events, the measured scatter in $N_{\\rm H}$ or \\tv\\ is very low, always $<$2. It seems then that assuming a single average value of \\tv\\ for all clouds in a given \\C\\ model is justified.    $\\bullet$ $Y$: the IR SED modeling and interferometry indicate the presence of dusty clouds out to $Y \\sim 20-25$ (e.g., RA11), although emission from beyond this radius may be difficult to separate from the surrounding host galaxy emission, especially for emission longer than $\\sim 10\\mic$. Our inferred locations of the X-ray-absorbing clouds cover out to $Y \\sim 20$, in agreement with the IR observations. It should be noted though that clouds at much larger distances from the black hole might exist, but would only be detectable with even longer sustained X-ray monitoring campaigns of a larger number of targets.   $\\bullet$ $\\sigma$, $i$ ($\\theta$) and \\No: we observe eclipse events in type I objects, and derive non-zero values of $P_{\\rm ecl}$. If it is the case that all type Is are oriented relatively face-on, then clumpy-torus models featuring sharp-edged tori are strongly disfavored, while models featuring a gradual decline of cloud number when moving away from the equatorial plane are favored (see left- and right-hand panels of Nenkova et al.\\ 2008b, Fig.~1, respectively). IR SED fitting likewise disfavors sharp-edged tori, which would produce a dichotomy in the SED shapes that is not observed (Nenkova et al.\\ 2008b).  We can derive our constraints on $\\sigma$ and \\No\\ assuming that an observed baseline level of absorption $N_{\\rm H,base}$ corresponds to an average of \\Ni\\ clouds along the line of sight ($\\Ni = 0$ in the case of sources normally lacking X-ray obscuration). We also assume that an eclipse denotes $\\Ni \\rightarrow \\Ni+1$. In the case of Cen~A, to repeat the exercise from Rothschild et al.\\ (2011) and Rivers et al.\\ (2011b), $\\Delta$$N_{\\rm H}$/$N_{\\rm H,base} \\sim 3$, suggesting $\\Ni \\sim 2$. From IR interferometry, Burtscher et al.\\ (2010) model a best-fitting value for the torus inclination of $i=63\\degr$, or $\\theta = 27\\degr$.  The median value for $\\sigma$ modeled by RA11 was $\\sigma = 20\\degr$; in order to agree with this value of $\\sigma$, one needs $\\No =12$.  In both cases of NGC~5506 and Mkn~348, $\\Delta$$N_{\\rm H}$/$N_{\\rm H,base} \\sim 2$, suggesting $\\Ni \\sim 1$. Due to the absence of external information on $i$, we can only derive constraints on the ratio $\\theta/\\sigma$. If we assume that $\\No = 10$, $\\theta/\\sigma$ = 1.52. That is, if such a system is inclined relatively edge-on ($\\theta\\rightarrow 0$), as is commonly suspected for type~IIs, then $\\sigma$ must be rather small, indicating a tightly flattened distribution. Otherwise, the inclination of the system would have to deviate more from an equatorial view. This exercise assumes that all obscuring material along the line of sight is in the form of clouds. If the baseline absorption $N_{\\rm H}$ is not due to clouds (see $\\S$5.6), then $\\Ni \\sim 0$, suggesting that the inclination of the system may not be close to edge-on, and/or that $\\sigma$ is relatively small in these systems. A similar condition, $\\Ni \\sim 0$, applies to sources normally lacking X-ray absorption, e.g., the type~I objects in our sample.   For a \\C\\ torus the probability of the AGN being obscured -- a single dust cloud along the line of sight is sufficient -- is given by \\begin{equation} \\label{eq:pobsc} \\pobsc(\\No, \\sig, i) = 1 - \\exp\\left\\{-\\No \\cdot \\exp\\left\\{ -\\left(\\frac{90-i}{\\sigma}\\right)^{\\!\\!\\!2}\\right\\}\\right\\} \\end{equation} (e.g., Nenkova et al.\\ 2008b; Nikutta et al.\\ 2009).  The argument to the outer exponential is the (negated) number of clouds along the line of sight \\Ni, thus $\\pobsc = 1 - \\exp\\{-\\Ni\\}$.  We can attempt to constrain the \\C\\ parameters \\No, \\sig, and $i$ by identifying ranges of \\pobsc\\ consistent with our X-ray monitoring data. In Fig.~\\ref{fig:pobscured} we plot \\pobsc\\ as a function of 2 of the 3 parameters, with each row of panels sampling the third parameter. The scale is linear (see color bar).  In each panel we overplot as contour lines $\\overline{P_{\\rm ecl}}$, the inferred values of the probability to catch an AGN while being obscured by a torus cloud, as computed in $\\S$\\ref{sec:sect43} (Table~\\ref{tab:finalprobs}).  The solid lines correspond to the minimum and maximum values of $\\overline{P_{\\rm ecl}}$, and dotted lines are best-estimate values. The white lines thus outline, given the data and the selection function, the range of permitted values of \\pobsc\\ for type I objects, and the black lines trace the allowed range of probabilities for type II objects. The dotted lines are the best-estimate values of \\pobsc.  \\begin{figure*} \\center \\includegraphics[width=0.8\\textwidth]{fig11.eps} % \\caption{\\C\\ parameter constraints based on the theoretical probability \\pobsc\\ to see the AGN obscured by torus clouds. The obscuration probability is rendered as a color scale in each panel (see color bar). The top panels show \\pobsc\\ as a function of torus angular width $\\sigma$ and viewing angle $i$ for four different values of $\\No$. The other two combinations of the three parameters that enter \\pobsc\\ are depicted in the middle and bottom rows. Overplotted as contour lines in each panel are the inferred ranges of $\\overline{P_{\\rm ecl}}$ that are compatible with our X-ray monitoring data (see Table~\\ref{tab:finalprobs}). White and black solid lines outline the compatible ranges for type I and II objects, respectively. The dotted lines indicate best estimates in both cases.} \\label{fig:pobscured} \\end{figure*}  These inferred minimal and maximal probabilities outline the parameter ranges permissible by our extensive monitoring programmes. Given the observed (and not observed!) eclipsing events, any torus parameters which lie outside the enclosed ranges are unlikely to be found in real sources. We remind the reader, however, that the probability values $\\overline{P_{\\rm ecl}}$ we derived in $\\S$4 refer only to lowly-ionized/neutral, full-covering clouds with $N_{\\rm H} \\ga 10^{22}$ up to $\\sim10^{25}$~cm$^{-2}$.  When eclipses are due to clouds with higher ionizations and/or lower column densities, and partial-covering events, the values will be higher.  Since three parameters (\\No, \\sig, $i$) affect \\pobsc\\ with some mutual dependence on each other, it is not straightforward to give one-dimensional ranges of excluded parameter values. However, for physically reasonable values some constraints can be given ad hoc. For instance, if \\No\\ is between 5 and 15 (see, e.g., Nenkova et al.\\ 2008b), it is clear from Fig.~\\ref{fig:pobscured} (bottom row) that \\sig\\ must be smaller than $\\approx 45\\deg$ (even for type II sources); only at the most pole-on viewing angles, \\sig\\ can be slightly higher. Similarly, if, for instance, \\sig\\ is secured by external information, for the reasonable range of \\No\\ the permitted viewing angles $i$ are very well-constrained (see middle row of Fig.~\\ref{fig:pobscured}). Not surprisingly, the permissible parameter values in type~II sources are on average higher than in type~I sources, because \\pobsc\\ \\emph{increases} when any of the three (\\No, \\sig, or $i$) grows. There is some overlap, however, and is most likely indicative of \"intermediate\" types of AGN.  Overall, Fig.~\\ref{fig:pobscured} shows that whenever one of the three parameters can be secured by other means, the other two allow only relatively narrow ranges. In practice, external information is sometimes available, for instance on the orientation of the AGN system from radio jets or from narrow-line ionization cones, and can be used to infer the two other parameters. We are not attempting a full Bayesian parameter inference on the torus parameters here. Suffice it to say here that our results as depicted in Fig.~\\ref{fig:pobscured} translate into information on the \\emph{priors} $P(\\bm\\theta)$ of the Bayesian inference problem $P(\\bm\\theta|D) \\propto P(\\bm\\theta) \\cdot P(D|\\bm\\theta)$. Here $\\bm\\theta = (\\No, \\sig, i, \\ldots)$ is the vector of model parameters to be estimated, $P(D|\\bm\\theta)$ is the \\emph{likelihood} that a set of parameter values generates a model compatible with the data, and $P(\\bm\\theta|D)$ is the full parameter \\emph{posterior} (also called inverse probability). Faced with perfect ignorance of the distribution of parameter values before modeling, uniform priors are often assumed. If there is independent knowledge available on some parameters it is imperative to incorporate this information into the Bayesian priors. For example, if the orientation of an AGN was known to be pole-on, one could assume a narrow Gaussian shape on the prior of $i$. We presume that then the Bayesian inference process would yield parameter intervals for \\No\\ and \\sig\\ commensurate with the lower-left panel of Fig.~\\ref{fig:pobscured}.   The \\pobsc\\ function transitions quite steeply from very low to very high probabilities. Our derived probability ranges (see Table~\\ref{tab:finalprobs}) typically bracket most of the transition region, and also exclude a large fraction of the parameter space that corresponds to extremely low or extremely high \\pobsc. We presume that a source such as NGC~1365, aligned so that \\Ni\\ (average number of clouds along the line of sight) is $\\approx 1$, would be just at the high end, or slightly above, of the type II region encompassed by the black lines in Fig.~\\ref{fig:pobscured}.      \\subsection{Can clumpy-torus models be applied to type II objects?}  %  As detailed in Appendix~A and summarized in $\\S$3, \\textit{RXTE} provided sustained long-term monitoring for durations $\\geq$0.6~yr for nine type II objects. In eight of them, $HR1$ is consistent with remaining constant for durations $\\geq$0.6~yr (we include the 2011 monitoring of Mkn~348 here). This constancy comes despite variations of typical factors $\\sim2-4$ in 2--10~keV continuum flux. For example, excluding the $\\la$1-d event in NGC~5506, there are no strong variations in $HR1$ (sustained trends above 2$\\sigma$ and/or $\\sim1.5 \\times \\langle HR1 \\rangle$) lasting longer than 1~d for a period of 8.39~yr.  Each of these sources routinely show evidence for the presence of absorbing gas with $N_{\\rm H} \\sim 10^{21-23}$~cm$^{-2}$ in their X-ray spectra, suggesting a baseline level of absorption $N_{\\rm H,base}$ that is consistent with being constant.  Using $HR1$ and its distribution, and assuming for simplicity that $\\Gamma$ remains constant for each object throughout the monitoring, we estimate the maximum value of $\\Delta$$N_{\\rm H}$ possible without significantly varying $HR1$. Those values are listed in Table~\\ref{tab:boringsy2s}, and for the brightest and best-monitored objects is typically $\\sim1-4\\times10^{22}$~cm$^{-2}$.  REN02 claim typically 20--80 per cent variations in $N_{\\rm H}$ in most of their sample of type II objects, using individual observations from multiple X-ray missions spanning $\\sim$25~yr.  In this paper, we rely on measurements from one instrument only, and so we do not face systematic uncertainties associated with cross-calibration between various X-ray missions or with fitting spectra data over different bandpasses.  For our eight type IIs, we can rule out variations above $\\sim20-30$ per cent for IRAS~04575--7537 and Mkn~348 (2011 monitoring only), and variations above $\\sim50-60$ per cent for NGC~1052, NGC~4258, and NGC~5506.   In the context of the clumpy-torus paradigm, having a non-zero value $N_{\\rm H,base}$ requires a non-zero number of clouds along the line of sight at all times, with the observed value of $N_{\\rm H,base}$ being the sum of the values of $N_{\\rm H}$ for the individual clouds. Assuming that all clouds have the same physical characteristics ($n$, $N_{\\rm H}$, etc.), the measured long-term constancy in $N_{\\rm H,base}$ would require that the number of clouds along the line of sight $\\Ni$ either remain constant for many years or is sufficiently large such that the addition/subtraction of one cloud intersecting or transiting along the line of sight does not change the total measured value of $N_{\\rm H}$ by a detectable amount.  However, with most clouds modeled with $\\C$ having $\\tau_{\\rm V} \\sim 30-100$, which corresponds to $N_{\\rm H} \\sim 6-20 \\times 10^{22}$~cm$^{-2}$ assuming the Galactic dust/gas ratio, values of $N_{\\rm H,base}$ of $10^{22}-10^{23}$ cm$^{-2}$ imply that $\\Ni$ cannot be more than a few.   We can explore if it is feasible to have the observed duration for each source correspond to a period where $\\Ni$ remains constant.  That is, we assume that all cloud contributing to the total observed value of $N_{\\rm H}$ remain in the line of sight for the duration, and we assume for simplicity that every cloud has identical $N_{\\rm H}$.  We also assume that every such cloud is $D_{\\rm cl}$ = 1 light-day in diameter and on Keplerian orbits.  Furthermore, we assume that the fastest velocity (closest to the black hole) cloud has just started to transit the line of sight when the long-term monitoring began, and that the cloud begins to leave the line of sight just when the monitoring ends.  In the case of NGC~5506 ($M_{\\rm BH}=10^{7.94} \\Msun$), an 8.39-year-long transit by such a cloud would require it to be located at least $r_{\\rm cl} = G M_{\\rm BH} (t_{\\rm d}/D_{\\rm cl})^2$ = 39 pc from the black hole.  For NGC~1052 ($M_{\\rm BH}=10^{8.19} \\Msun$; Woo \\& Urry 2002), a 4.56-year transit would place it $\\geq20$ pc from the black hole.  For NGC~4258, ($t_{\\rm d}$=6.81~yr, 2005 March -- 2011 December; $M_{\\rm BH} = 3.9\\times 10^7 \\Msun$, Miyoshi et al.\\ 1995; Herrnstein et al.\\ 1999), $r_{\\rm cl} \\geq$ 11 pc.\\footnote{Resulting values of $r_{\\rm cl }$ for NGC~2992 ($t_{\\rm d}$=0.90~yr; $M_{\\rm BH} = 5.2\\times10^7~\\Msun$, Woo \\& Urry 2002), Mkn~348 ($t_{\\rm d}$=0.98~yr, 2011 only), and NGC~7314 ($t_{\\rm d}$=1.56~yr, 1999--2000; $M_{\\rm BH} = 1.4 \\times 10^6$~$\\Msun$, Vasudevan et al.\\ 2010) are 0.26, 0.09, and 0.02 pc, respectively.} A cloud diameter of 0.1 light-days, also a feasible value given the range listed in Table~\\ref{tab:clumpsizemass}, would yield values of $r_{\\rm cl}$ an order of magnitude higher.  The clumpy-torus models typically assume that clouds cannot exist in large numbers out to several tens of $r_{\\rm d}$, or very roughly 10, 5 and 0.5 pc for NGC~5506, NGC~1052, and NGC~4258, respectively\\footnote{Bolometric luminosities for NGC~5506, NGC~1052 and NGC~4258 from Vasudevan et al.\\ (2010), Woo \\& Urry (2002), and Lasota et al.\\ (1996), respectively.}.  In these objects, it is thus not very likely that the same clouds used in the $\\C$ models can also be responsible for the non-variable baseline level of absorption.  Furthermore, Lamer et al.\\ (2003) and Rivers et al.\\ (2011b) demonstrated that the transiting clouds responsible for the NGC~3227/2000--1 and Cen~A/2010--1 events were non-uniform, with a number density increasing towards the center.  If it were the case that these two clouds were representative of all AGN clouds in terms of their density profiles, then any single cloud moving across the line of sight should produce a secular variation in observed $N_{\\rm H}$ and $HR1$. However, we find no evidence for any strong secular trend in any type II object, to the limits indicated by Table~\\ref{tab:boringsy2s}.  In addition, a transiting spherical cloud with a constant \\textit{column} density profile would require $n$ to decrease as one goes from the cloud edge to the center.   We conclude that it is difficult for clumpy-torus models to satisfactorily explain the constancy of $N_{\\rm H,base}$, especially in the cases of the few longest-monitored and X-ray brightest type II objects observed with \\textit{RXTE}.  It is thus more likely that $N_{\\rm H,base}$ arises in a non-clumpy, highly-homogeneous (to within the above $\\Delta$$N_{\\rm H}$ limits) medium; such a medium could be located at any radial location along the line of sight, subject to restrictions from the medium's ionization parameter.  That is, the cloud responsible for the $\\la1$-d eclipse in NGC~5506 and the material responsible for $N_{\\rm H,base} \\sim 2\\times10^{22}$~cm$^{-2}$, for instance, are likely physically separate entities.  One possibility to explain $N_{\\rm H,base}$ is a smooth, relatively low-density, dusty intercloud medium in which the relatively higher density clouds are embedded (Stalevski et al.\\ 2012).  To be consistent with the approximate range of baseline values of $N_{\\rm H}$ typically measured in type IIs' X-ray spectra ($0.3 - 30 \\times 10^{22}$~cm$^{-2}$), the intercloud medium could have values of $\\tau_{\\rm V} \\sim 1.6 - 160$ or $\\tau_{10\\micron} \\sim 0.07 - 7$.\\footnote{We employ the optical properties of the composite silicate/graphite grains used in the \\C\\ models, with $\\tau_{\\rm V} = 23.6\\tau_{10\\micron}$.}  However, we do not see evidence for such a medium in type~I AGN.  Another possibility we explore is X-ray-absorbing gas distant from the black hole and associated with the host galaxy, as has been suggested e.g., by Bianchi \\etal\\ (2009) for NGC~7582. For example, dust lanes/patches associated with the host galaxy are considered a candidate for the obscuration of BLR lines in some systems (Malkan et al.\\ 1998).  With the exception of Cen~A (see $\\S$\\ref{sec:cenasmalldip}), we rely on Shao et al.\\ (2007) and Driver et al.\\ (2007), who derive the inclination dependence of dust extinction for disk-dominated galaxies. In the optical bands, the extinction is typically $\\la1-2$ mag, derived over a range of inclination angles from face-on to nearly edge-on ($i \\sim 80-85 \\degr$), corresponding to a maximum $N_{\\rm H} \\sim 4 \\times 10^{21}$~cm$^{-2}$ for a nearly edge-on disk. If the obscuring dust is distributed uniformly throughout the disk of the host galaxy, then such a component may explain relatively low observed values of $N_{\\rm H,base}$ in, for instance, NGC~2992 ($N_{\\rm H,base} \\sim 4 \\times 10^{21}$~cm$^{-2}$) with the disk component of its host galaxy oriented close to edge-on (Jarrett et al.\\ 2003). However, for objects with host disk inclinations far from edge-on (e.g., NGC~4258 and Mkn~348: Hunt et al.\\ 1999; Jarrett et al.\\ 2003) and/or for objects with $N_{\\rm H,base} \\sim$ a few $\\times 10^{22}$~cm$^{-2}$ and higher, it is difficult for such a distribution of dust to be associated with $N_{\\rm H,base}$. Of course, we cannot rule out the possibility of having more concentrated regions of gas lying along the line of sight not associated with the large-scale disk structure (such as a Giant Molecular Cloud). Dust-\\textit{free} X-ray-absorbing gas in any of these configurations is also a candidate.      A final possible explanation for constancy in $N_{\\rm H}$ can be, as mentioned above, the presence of a large number of low-density clouds, with $\\Ni$ remaining nearly constant.  However, this would require each cloud to have $N_{\\rm H} \\ll 10^{22}$~cm$^{-2}$.  Such values contrast with those used by Nenkova et al.\\ (2008b), RA09, RA11, and the values measured in this paper, but such clouds are not physically implausible.  Contributions from such clouds, if they are dusty, to the total IR emission are negligible. However, the sum of many clouds along the line of sight that each have $N_{\\rm H} \\ll 10^{22}$~cm$^{-2}$, be they dusty or dust-free, could yield appreciable X-ray absorption.   \\subsubsection{A small dip in $N_{\\rm H,base}$ in Cen~A} % \\label{sec:cenasmalldip}  Up to this point, our discussion of the baseline level of X-ray obscuration in type IIs has centered on constant (to within our sensitivity limits) values of $N_{\\rm H, base}$. However, $N_{\\rm H, base}$ for Cen~A warrants more attention in light of the small dip in early 2010.  The famous dust lane crossing the host elliptical of Cen~A supplies 3--6 mag.\\ of optical extinction (Ebneter \\& Balick 1983), corresponding to $\\sim5-10 \\times 10^{21}$~cm$^{-2}$, far short of observed values of $N_{\\rm H,base} \\sim 1-2 \\times 10^{23}$~cm$^{-2}$.  It is generally accepted that a more compact and higher density distribution is responsible for observed values of $N_{\\rm H}$ in X-ray spectra of Cen~A.  One possibility is that we have witnessed one cloud leaving the line of sight followed by another cloud entering the line of sight two months later. Best-fitting values of $N_{\\rm H}$($t$) indicate a drop from $21.7\\pm0.9$ to $18.6^{+0.9}_{-0.8}$ and returning to $N_{\\rm H}$($t$) $ = 21.9 \\pm 0.7 \\times 10^{22}$~cm$^{-2}$. These values suggest a scenario in which $\\Ni$ dropped temporarily from $\\sim$ 7 to 6, with each cloud contributing $\\Delta$$N_{\\rm H} \\sim 3 \\pm 1 \\times 10^{22}$ cm$^{-2}$, or less than half the column inferred to exist for the clouds causing the 2003--2004 and 2010--2011 \\textit{increases} in total $N_{\\rm H}$. If all these clouds are part of the same clumpy structure and have a common origin, then the assumption of each cloud in a given AGN having exactly uniform values of $N_{\\rm H}$ is an oversimplification. $N_{\\rm H} = 3 \\pm 1 \\times 10^{22}$~cm$^{-2}$ is not an implausible value for a cloud, and only a factor of 2--3 less than that found for other clouds detected in our sample.   $N_{\\rm H,base}$ in Cen~A could also be attributed to non-clumpy, spatially-extended material that is not entirely homogeneous; perhaps this component can be associated with an intercloud medium. Here, $N_{\\rm H,base}$ is usually $\\sim 21.8 \\times 10^{22}$~cm$^{-2}$, but there exists an \\textit{underdense} region with $\\Delta$$N_{\\rm H} = -3 \\times 10^{22}$~cm$^{-2}$ that transited the line of sight in 2010. Assuming an arbitrary distance to the black hole of 100--200 light-days (the location of the absorbing clouds as per $\\S$5.2), the $\\sim$80-d duration of the dip implies that the underdense region is on the order of half a light-day across.   Longer-term variability in $N_{\\rm H,base}$ may also exist: Rothschild et al.\\ (2011) measured values of $14-19\\times10^{22}$~cm$^{-2}$ during the 1996--2009 \\textit{RXTE} observations, excluding the 2003--2004 points. Hopefully, future X-ray monitoring will further resolve the variability in $N_{\\rm H, base}$ and allow us to discriminate among the above scenarios."
    },
    "1402/1402.7195_arXiv.txt": {
        "abstract": "We review metrics to assess the radio quietness of sites used for radio astronomy.  Concise metrics are needed to compare candidate sites for new telescopes, to monitor the quality of existing sites, and to design telescopes to work well at a given site.  Key points of assessment are the receiver dynamic range required for the strongest interferers and the expected fraction of spectrum available for sensitive astronomical measurements.  We propose three metrics: (1) total radio frequency interference power, (2) interference-to-noise power ratio and (3) time-frequency occupancy.  Box plots of these metrics summarise large quantities of information, highlight expected ranges of interfering signal properties, and aid comparisons of sites and other factors of interest.  We provide examples for Square Kilometre Array phase one deployment in Australia based on measurements made for the selection of this site.  The Square Kilometre Array will be the largest radio telescope in the world. ",
        "introduction": "Our goal is to assess the impact of man-made radio frequency interference (RFI) on observatory viability and efficiency via as few numbers and plots as possible.  Carefully presented summary metrics can aid efficient and fair comparison of candidate telescope sites and identify important trends at existing sites.  RFI surveys produce very large data sets that prohibit inspection of every spectrum.  Studying individual spectra may be necessary to understand fine details.  However, a high-level metric based analysis should be made first to identify where such detailed effort is best spent.  Concise metrics are required to summarise the effects of RFI on (1) the viability of telescope receiver hardware and (2) availability of spectrum for sensitive astronomy.  Availability of spectrum for sensitive astronomical measurements is the key issue for scientific efficiency of an observatory.  However, this should not be considered independently of the viability and cost of RFI-robust receivers and signal processing.  We propose a high-level analysis based on summary statistics of total RFI power, interference-to-noise power ratio (INR), and time-frequency occupancy.  Statistics of the first two metrics summarise the range of RFI powers that receivers must cope with at a given site.  Good results for these are required so that a practical and economical telescope can be built to realise the otherwise theoretical spectrum availability indicated by the third metric.  \\emph{Total RFI Power} statistics help engineers design analog receiver front ends with sufficient headroom.  These front ends must operate linearly for the majority of expected RFI power levels.  Box plots of total RFI power like Fig. \\ref{fig:boxplot} neatly summarise the typical range of total RFI power and additionally highlight the strongest powers that may damage a receiver.  Engineers can quickly gauge whether a design has sufficient headroom by adding their own power limit lines for receiver damage and input 1~dB compression to such plots.  \\begin{figure}[b!] \\centerline{\\psfig{figure=boxplot-au-all.eps,height=64.54mm} } \\caption{Box plot summarising total RFI power detected in SKA Phase I bands at the Australian SKA site.  Boxes extend from the $25^{\\text{th}}$ to $75^{\\text{th}}$ percentiles and are broken by a horizontal line at the median.  Whiskers extend to 10$^{\\text{th}}$ and 90$^{\\text{th}}$ percentiles, and dots mark maxima and minima.  These statistics are calculated over estimates of total RFI power from each of the 853 usable 60~s integration spectra.  Thickened lines show expected receiver system noise power $P_{sys} = k\\overline{T}_{sys}B$ in each band.   } \\label{fig:boxplot} \\end{figure}  \\emph{Interference-to-Noise Power Ratio} (INR) statistics help engineers design digital back ends with sufficient headroom.  INR is the ratio of total power in detected RFI signals to total noise power expected from an astronomy receiver in a specified band.  Typical range and and extreme values of INR can be read from Fig. \\ref{fig:boxplot} by comparing the box plots to the thickened system noise power lines.  This helps engineers assess the number of bits required to keep digital errors below specified thresholds.  \\emph{Time-Frequency Occupancy} statistics help astronomers estimate the overall efficiency of an observatory.  They indicate what fraction of data is affected by RFI at a particular sensitivity threshold.  This fraction of data cannot be used for science without invoking additional signal processing for RFI mitigation.  Fig. \\ref{fig:occupancy} shows an example.  \\begin{figure}[t!] \\centerline{\\psfig{figure=occupancy-time-freq.eps,height=64.54mm} } \\caption{Time-frequency occupancy in 100~MHz bands at the Australian SKA site.  This is the percentage of 27.4~kHz$\\ \\times \\ $60~s data points affected by RFI in each 100~MHz$\\ \\times \\ $2~hr high-sensitivity spectrogram.  Variations over all data from 8 spectrograms covering 4 directions $\\ \\times \\ $ 2 polarisations are shown by boxes extending from the 25$^{\\text{th}}$ to 75$^{\\text{th}}$ percentiles that are split by horizontal lines at median values.  Whiskers extend to 10$^{\\text{th}}$ and 90$^{\\text{th}}$ percentiles.  The dashed line shows mean time-frequency occupancy for all data.} \\label{fig:occupancy} \\end{figure} ",
        "conclusions": "RFI metrics to guide practical receiver design are as important as metrics estimating available spectrum for astronomy assuming an ideal receiver.  These issues should not be separated.  We hope this work inspires an open discussion  of these and other metrics.  Progressing this work to a standard procedure for high-level analysis and communication of RFI data would significantly benefit the international astronomical community and the large international projects it is pursuing.  These involve many stakeholders that must share information across many disciplines to select good sites, protect good sites, and design hardware that works well at these sites.  Further measurements at both SKA sites are required to collect the information required to design successful SKA receivers.  These measurements need to be made with higher dynamic range receivers, peak detection, and triggered recording of transients.  These, along with further high-sensitivity spectrogram measurements, should be carried out over long periods to better understand variations in spectrum usage and propagation over time-scales from nanoseconds to years."
    },
    "1402/1402.5968_arXiv.txt": {
        "abstract": "{Nova Delphini 2013 was identified on the 14th of August 2013 and eventually rose to be a naked eye object.} {We sought  to study the behaviour of the object in the run-up to outburst and to compare it to the pre-outburst photometric characteristics of other novae.} {We searched the Pan-STARRS\\,1 datastore to identify pre-outburst photometry of Nova Del 2013 and identified twenty-four observations in the 1.2 years before outburst.} {The progenitor of Nova Delphini showed variability of a few tenths of a magnitude but did not brighten significantly in comparison with archival plate photometry. We also found that the object did not vary significantly on the approximately half hour timescale between pairs of Pan-STARRS\\,1 observations.} {} ",
        "introduction": "Nova Delphini 2013 (=V339 Del) was first identified as a bright 6.8 magnitude source by Koichi Itagaki on the 14th of August 2013 and announced in AAVSO Alert Notice 489 \\citep{Waagen2013}. Since then this object has been extensively monitored using radio continuum \\citep{Roy2013}, near-infrared \\citep{Stringfellow2013}, optical \\citep{Darnley2013,Tarasova2013a}, X-ray (with \\citealt{Page2013} first detecting emission 36 days after discovery) and gamma-ray \\citep{Hays2013} observations. Initially the nova exhibited strong P~Cygni profiles in the Balmer and He~I lines \\citep{Tarasova2013a}, these profiles disappeared from the hydrogen lines on the 19th of August \\citep{Darnley2013} but remained for other species (He~I, O~I, Fe~II).  The progenitor of this nova has been identified as USNO-B1.0~1107-0509795 which had an unfiltered magnitude of 17.1 on the 14th of May 2013 \\citep{Denisenko2013}. \\cite{Wren2013} report a pre-discovery detection as the nova was brightening. Data on this object prior to outburst remain relatively scarce. Denisenko reports a counterpart in the UV from GALEX while there are a handful of detections in POSS photographic plate catalogues. \\cite{Tang2013} report three pairs of H$\\alpha$ on/off measurements but find no strong evidence that the progenitor is an H$\\alpha$ emitter. \\cite{Munari2013} identified the progenitor in Asiago photographic plate data, finding 12 detections and 13 upper limits on the object's brightness. Over the course of three years of observations from 1979 to 1982, the progenitor showed variation in the $B$-band with a total amplitude of 0.9\\,mag. \\cite{Munari2013} also stacked three observations taken by the APASS survey \\footnote{http://www.aavso.org/apass} in April 2012 to provide additional $B$ and $V$ photometry.  \\subsection{Pre-outburst observations of novae} \\cite{Robinson1975} presented a landmark study of pre- and post- outburst data for novae from archival photographic plates. Out of 18 targets with both pre- and post-outburst magnitude measurements, he identified only one object (BT Monocerotis) with significantly different brightnesses.  Out of 12 targets with well-sampled pre-outburst light curves, five showed an apparent increase in luminosity in the few years before outburst, and two showed significantly increased variability in the lead-up to their outbursts.  These claims were later re-examined by \\cite{Collazzi2009}, who combined subsequent remeasurements of plate data both from their own work and from other sources (whereas \\citealt{Robinson1975} relied on the available pool of published photometry from the plates, which can suffer from systematic calibration errors up to $\\gtrsim1$~mag). They concluded that BT Mon actually showed no significant change between its pre- and post-outburst brightness and did not reproduce four of the five claims of pre-outburst brightening. They confirmed that the one remaining object from the \\cite{Robinson1975} pre-outburst brightening sample (V533 Herculis) did, in fact, increase in brightness by $\\sim1.3$~mag in the $\\sim1.5$~years leading up to outburst.   \\cite{Collazzi2009} identified an additional nova (V1500 Cygni) that showed pre-outburst brightening, rising nearly 7~mag in the month before its outburst\\footnote{Although post-outburst behavior is not explicitly the topic of this paper, we note that V1500 Cygni was also one of the five (out of 30) novae that  \\cite{Collazzi2009} found to have brightened by more than a factor of 10 after outburst, while the remaining majority of novae had essentially no change in quiescent brightness across the outburst.}.  In addition,  \\cite{Collazzi2009} found no case for a class of novae with a significant increase in variability leading up to the outburst, noting that one such object claimed by \\cite{Robinson1975} (RR Telescopii) has since been found to be a symbiotic star and that the other (V446 Herculis) actually did not exhibit any change in its variability behavior before and after outburst. They also note that changes in variability amplitude in general are characteristics of novae in quiescence and may not be associated with the outburst -- they conclude that none of the novae they examined showed a significant change in variability that could be linked to the outburst. It should be noted that \\cite{Collazzi2009} also found also found a pre-outburst dip in the recurrent nova T CrB. Recent additional archival photometry unearthed by \\cite{Schaefer2014} indicates that this object was $\\sim1\\,mag$ brighter in the 8 years leading up to this brief pre-erruption dip. \\cite{Schaefer2013} also report a brief, temporary brightening of $\\sim$1\\,mag of  another recurrent nova, T Pyx, prior to outburst.  The total brightness of cataclysmic variables (CVs) with accretion disks -- of which classical novae are a subset -- are dominated by accretion-generated luminosity over a wide range of the electromagnetic spectrum (from the near-IR to the far-UV).  As such, variability in the long-term\\footnote{``Long-term'', in this case, is relative since it might easily be measured in units of the CV's orbital period.  At typical values of $\\sim1$--$10$~hr, even many orbital cycles span an interval that is, in an absolute sense, not very long.  CVs, of course, famously display a wide range of short timescale, low amplitude variability - ``flickering'' -- but this is likely related to the effects of viscous turbulence in the disk and/or accretion stream and not to actual changes in the mean mass transfer rate from the donor star \\citep{Scaringi2013}.}. mean brightness is directly proportional to changes in the mass transfer rate from the low mass donor star, through the disk, and onto the white dwarf (WD) primary star.  The origin of a classical nova outburst is a thermonuclear runaway in the thin layer of accreted hydrogen-rich matter on the surface of the WD.  The total mass in this layer required to trigger the runaway is a few $\\times10^{-3}$~M$_{\\odot}$ for a WD with $M_{\\rm WD}=0.6$~M$_{\\odot}$ (i.e., the average WD mass) to a few $\\times10^{-5}$~M$_{\\odot}$ for WDs approaching the Chandrasekhar mass (see \\citealt{Warner2003} and references therein).  Even under the optimistic assumption of a persistent high CV mass transfer rate of $\\sim10^{-9}$~M$_{\\odot}$~yr$^{-1}$, the timescale to grow this accreted envelope to the critical mass is $\\gtrsim10^{4}$--$10^{6}$~yr.  This is the root of the reason that \\cite{Collazzi2009} found Robinson's conclusion that almost 50\\% of novae show a pre-outburst brightness increase to be ``unsettling''.  They discuss in some detail the arguments against a corresponding pre-outburst increase in the mass transfer rate being the ``last straw'' that triggers the nova outburst.  In summary, (a) the amount of extra mass accumulated during the rise in brightness is miniscule compared to the total mass required to trigger the outburst, and (b) the energy flux carried by the extra accreted matter produces an insignificant increase in the temperature at the base of the accumulated layer on the WD.  The result of  \\cite{Collazzi2009}, two out of 22 novae ($<10$\\%) displaying a pre-outburst rise in brightness, might well be less unsettling, but it is still problematic:\\ as they note, the close proximity to the nova outbursts in both V533 Her and V1500 Cyg suggests that the pre-outburst brightness increases are somehow causally connected to the outbursts, although they cannot have directly triggered the outbursts.  In the majority of classical novae, it appears, the outburst should arrive as a surprise, without previous announcement in the light curve.  It is the currently unexplained exceptions to this situation, like V533 Her and V1500 Cyg, that might offer clues that will help to further illuminate the nova outburst mechanism.  Similarly, although both \\cite{Robinson1975} and  \\cite{Collazzi2009} agree that the majority of novae have the same pre- and post-outburst mean quiescent brightness (implying that the outburst does not significantly alter the mass transfer process in the CV), the latter found that $\\sim15$\\% of novae {\\em do} show a post-outburst increase in brightness of $>2.5$~mag.  Again, the few exceptions to the majority behavior might offer the best insight into the nova outburst mechanism.  Obtaining post-outburst data on novae is easy, but for pre-outburst data, we must typically rely on serendipitous observations.  In this paper we describe pre-outburst photometry from Pan-STARRS\\,1 for Nova Del 2013. ",
        "conclusions": "We find that prior to its eruption, Nova Del 2013 did not show a significant change in photometric behaviour. Similar to most of the novae studied by \\cite{Collazzi2009}, the progenitor did not show a significant increase in brightness in the few years prior to outburst. The lack of significant variability between the components of Pan-STARRS\\,1 observation pairs indicates that the progenitor was not significantly variable on the timescale of approximately half an hour.  During the next several years, a similar amount of pre-outburst photometry will be available for the majority of northern novae from the ongoing Pan-STARRS project.  A number of similar ongoing and planned wide-field, temporal surveys will provide additional multi-epoch photometric coverage of nova progenitors over, collectively, a large fraction of the entire sky in both hemispheres (e.g., the Catalina Real-time Transient Survey, \\citealt{Drake2012}; the Palomar Transient Factory, \\citealt{Rau2009}; SkyMapper, \\citealt{Keller2013}).  Moving into the next decade, the more densely sampled observing cadence of the Large Synoptic Survey Telescope (LSST; \\citealt{Ivezic2008}) will provide even more data that can be used to characterize the behavior of novae during the (potentially long-term) run-up to outburst.  The arduous process of combing through archival plates in the hopes of serendipitously locating nova progenitors, and then processing those images to provide usable photometric data, will be increasingly supplanted by an ever-growing spatial and temporal database of accurate, homogeneous, digital photometry extending to very faint objects.  This will finally allow us to systematically assess the presence and characteristics of photometric changes displayed by classical novae in relation to their outbursts, providing crucial input and constraints to modelling the physics behind these accretion-induced thermonuclear events."
    },
    "1402/1402.0418_arXiv.txt": {
        "abstract": "In this paper we present a comprehensive study of the heavy quark production in ultra high energy cosmic ray interactions in the atmosphere considering that the primary cosmic ray can be either a photon, neutrino or a proton. The analysis is performed using a unified  framework -- the dipole formalism --- and  the  saturation effects, associated to the physical process of parton recombination, are taken into account. We demonstrate that the contribution of heavy quarks for cosmic ray interactions is in general non-negligible and can be dominant depending of the process considered. Moreover, our results indicate that new dynamical mechanisms  should be included in order to obtain reliable predictions  for the heavy quark production in $pp$ collisions at ultra high cosmic ray energies. ",
        "introduction": "Ultra high energy cosmic rays (UHECRs) remains a puzzle in physics. Although the existence of UHECRs with energies above $10^{11}$ GeV is now a well-established fact, the theoretical understanding of its origin and propagation is a subject of  strong interest and intense discussion \\cite{review_cosmic}.  In particular, the recent detection of  ultra-high energies (UHE) neutrinos by the IceCube Neutrino Observatory \\cite{Icecube} starts a new era in the neutrino physics. These and forthcoming data from IceCube and Pierre Auger Observatories may shed light on the Standard Model as well as  reveal aspects of new physics \\cite{review_neutrinos}. An important subject in cosmic ray physics is the flux of prompt leptons at the Earth which is related to   primary interactions at energies that can by far exceed the highest available in accelerators. This flux is directly associated with the charmed particle production and its decays, being  its estimation strongly dependent on the model used to calculate the charm production cross section and energy spectra \\cite{leptonflux}.  The quantification of the prompt leptons fluxes is essential, for instance, for neutrino physics, since the flux of prompt neutrinos provides the main background of the muon neutrinos from extra-galactic neutrino sources, which are being studied by the  IceCube Observatory.   In this paper we calculate the heavy quark production in UHECR interactions considering that the particle incident in the atmosphere can be either a photon, neutrino or a proton. As this interaction probes the theory of the strong interactions in a new kinematical range characterized by a center of mass energy of approximately 500 TeV, which is more than one order of magnitude larger than that of the Large Hadron Collider at CERN, we take into account the nonlinear corrections to the QCD dynamics effects associated to the gluon saturation physics \\cite{book_kov}. In this energy regime, perturbative Quantum Chromodynamics (pQCD) predicts that the small-$x$ gluons in a hadron wavefunction should form a Color Glass Condensate (CGC),  which  is characterized by the limitation on the maximum phase-space parton density that can be reached in the hadron wavefunction (parton saturation), with the transition being specified  by a typical scale, which is energy dependent and is called saturation scale $Q_{\\mathrm{sat}}$. Moreover, we calculate the heavy quark production using the color dipole approach, which gives a simple unified picture for this process in photon-proton, neutrino - proton and proton - proton interactions. Our goal is to  extend the  analysis performed in Ref. \\cite{victor_jhep} some years ago, where the main focus was to estimate the heavy quark cross section  in photon - air and proton - air interactions   taking into account nuclear effects. In particular, we extend the study for  neutrino - proton interactions and explicitly estimate, for the first time,  the  heavy quark contributions for the total $\\gamma p$, $\\nu p$ and $pp$ cross sections. Distinctly from Ref.  \\cite{victor_jhep}, the $\\gamma p$ and $pp$ interactions are analysed in detail, with the free parameters present in the color dipole predictions being constrained by the experimental data for the total $\\gamma p$ cross section and by the recent LHC data for the heavy quark production, which strongly reduces the uncertainty in  the predictions for ultra high cosmic ray energies. Moreover, the  heavy quark production in a double parton scattering process in cosmic ray $pp$ interactions is estimated for the first time.  It is important to emphasize that although there are another studies in the literature that calculate the heavy quark contribution for the total cross sections (See {\\it e.g.} Refs. \\cite{kms,reno,frankfurt,Marta_Rafal}), it is the first time that it is estimated using a unified approach to treat the distinct processes. In general, these studies have been performed for a specific process of the three  analysed here  considering different underlying assumptions, as {\\it e.g.} the order of the perturbative calculations,  the QCD evolution equations and/or choice of the parton distributions, heavy quark masses and factorization scale, which introduces a large theoretical uncertainty in the predictions. Distinctly of these previous studies, our predictions for the heavy quark contribution in $\\gamma p$, $\\nu p$ and $pp$ processes are based on a unified approach, considering a unique model for the QCD dynamics and for the value of the heavy quark masses.    The paper is organized as follows. In the next Section we present a brief review of the  color dipole picture, presenting the main formulas and assumptions present in the calculation of the heavy quark production in this approach. In Section \\ref{dynamics_section} we discuss the QCD dynamics at high energies and present the models used for the dipole - proton scattering amplitude used in our calculations. In Section \\ref{results_section} we present our results for the energy dependence of the heavy quark contributions for the  total  $\\gamma p$, $\\nu p$ and $pp$ cross sections. A comparison with previous predictions is presented and the magnitude of the gluon saturation effects is estimated. Moreover, the probability of the heavy quark production in a  double parton scattering process at ultra high cosmic ray $pp$ interactions is calculated.  Finally, in Section \\ref{summary_section} we summarize  our main conclusions. ",
        "conclusions": "\\label{summary_section}  Heavy quark production in hard collisions of hadrons, leptons, and photons has been considered as a clean test of perturbative QCD. This process provides not only many tests of perturbative QCD, but also some of the most important backgrounds to new physics processes,  which have motivated  comprehensive phenomenological studies carried out  at DESY-HERA, Tevatron and LHC. With the advent of the Pierre Auger and  IceCube observatories, the extension of these studies for ultra high energy cosmic ray interactions becomes fundamental. In this paper we estimate the heavy quark production at these energies  considering the primary particle as being either a photon, neutrino or a proton. For the first time, the heavy quark contribution for the total cross section  is estimated considering a unique framework -- the dipole formalism --- and taken into account saturation effects, associated to the physical process of parton recombination. Our results indicate that the contribution of heavy quarks is in general non-negligible and can be dominant depending of the process considered. In particular, for $pp$ collisions, the large value obtained for the charm contribution at cosmic ray energies can be interpreted as a signal of the breakdown of the factorization used to calculate the heavy quark cross section.  Moreover, they indicate that new dynamical mechanisms, as for instance, double parton scattering processes, should be included in order to obtain reliable predictions."
    },
    "1402/1402.0903_arXiv.txt": {
        "abstract": "We analyze the Aquarius simulations to characterize the shape of dark matter halos with peak circular velocity in the range $8<V_{\\rm max}<200$ $\\rm km/s$, and perform a convergence study using the various Aquarius resolution levels. For the converged objects, we determine the principal axis ($a \\geq b \\geq c$) of the normalized inertia tensor as a function of radius.  We find that the triaxiality of field halos is an increasing function of halo mass, so that the smallest halos in our sample are $\\sim 40 - 50\\%$ rounder than Milky Way-like objects at the radius where the circular velocity peaks, $r_{\\rm max}$. We find that the distribution of subhalo axis ratios is consistent with that of field halos of comparable $V_{\\rm max}$. Inner and outer contours within each object are well aligned, with the major axis preferentially pointing in the radial direction for subhalos closest to the center of their host halo. We also analyze the dynamical structure of subhalos likely to host luminous satellites comparable to the classical dwarf spheroidals in the Local Group. These halos have axis ratios that increase with radius, and which are mildly triaxial with $\\langle b/a \\rangle\\sim0.75$ and $\\langle c/a \\rangle\\sim0.60$ at $r \\sim 1$ kpc. Their velocity ellipsoid become strongly tangentially biased in the outskirts as a consequence of tidal stripping. ",
        "introduction": "\\label{sec:intro}  In the $\\Lambda$ cold dark matter cosmological paradigm structures build hierarchically, via the mergers of smaller objects \\citep{Press1974, Gott1975, White1978, Blumenthal1984}. As mergers proceed, the innermost regions of some of the progenitors survive, resulting in non-linear structures where a wealth of substructure orbits the center of an otherwise monolithic dark halo. Early $N$-body simulations showed that halos could host dozens of substructures, down to masses near the numerical resolution limit \\citep{Tormen1997, Tormen1998, Ghigna1998, Klypin1999b, Klypin1999, Moore1999}. For systems like the Milky Way, current numerical simulations have extended the dynamical range of resolved substructures by 4-5 orders of magnitude \\citep{Diemand2007, Diemand2008, Springel2008, Stadel2009}.  The properties of these substructures are of great interest since luminous satellites, such as the population of dwarf spheroidal (dSph) galaxies in the Local Group, are expected to be embedded in them \\citep{Stoehr2002, Strigari2007, Boylan2012, Vera2012}. Furthermore, the large mass-to-light ratios of dSph, which range from 10 s to 1000 s \\citep{Mateo1998, Gilmore2007, Walker2012}, indicate that their internal dynamics is dominated by the dark matter.  This suggests that the predictions of pure dark matter simulations may be directly confronted with observations of these systems. For instance, it has been suggested that they provide an optimal place to look for signals of dark matter self-annihilation processes \\citep{Kamionkowski2010} due to the natural enhancement in density and the lack of significant contamination from the baryonic component.   The availability of large samples of line-of-sight velocities for individual stars in dSph galaxies offers new tests of the predictions of $\\Lambda$ cold dark matter ($\\Lambda$CDM) regarding the structure of dark matter subhalos. For instance, studies of $N$-body numerical simulations have shown that the inner slope of the dark matter density profile is expected to be cuspy in CDM models \\citep{Navarro1996, Navarro1997}.  This seems to contrast with the somewhat shallower slopes and even constant density cores proposed to explain the motions of stars in local dwarf spheroidals \\citep{Amorisco2011, Walker2011, Amorisco2012, Jardel2012}. This, however, is a subject of active debate, since various authors have shown that the stellar kinematics of Milky Way dwarfs are also consistent with the NFW cuspy profiles \\citep{Battaglia2008, Walker2009, Strigari2010, Breddels2012,Breddels2013}.   Most of the dynamical modeling performed in the studies of Local Group dwarf spheroidals relies on simple assumptions about the structure of their dark matter component.  In particular, spherical symmetry and specific anisotropy profiles have been extensively assumed.  The orbital anisotropy has been taken to be constant \\citep{Richstone1986, Lokas2002, Lokas2005, Lokas2009, Walker2009}, or radially dependent \\citep{Kleyna2001, Wilkinson2002, Battaglia2008, Strigari2008, Wolf2010, Amorisco2011}, while in Schwarzschild modeling or in made-to-measure $N$-body methods it does not need to be assumed \\citep{Long2010,Jardel2012, Breddels2012}. The selection of geometric shape for the dark matter potential can also be relaxed. For example, \\cite{Hayashi2012} considered axisymmetric dark matter halos to model some of the Milky Way dSph galaxies.  For isolated galaxies, numerical experiments of $\\Lambda$CDM have clear predictions for these quantities. The shapes of (isolated) dark matter halos in the mass range $10^{10}-10^{15}$ M$_{\\odot}$, are generally found to be triaxial, with axis ratios depending on the mass of the object \\citep{Frenk1988, Dubinski1991, Warren1992, Cole1996, Thomas1998, Jing2002, Bailin2005, Kasun2005, Hopkins2005, Allgood2006, Bett2007, Hayashi2007, Kuhlen2007, Stadel2009, Diemand2009, Munoz2011, Schneider2012}. In terms of their internal kinematics, the velocity ellipsoid is close to isotropic near the center of halos and becomes mildly radial towards the outskirts \\citep{Wojtak2005,Hansen2006, Ludlow2011}.   For subhalos, however, less is known because of the demanding numerical resolution needed to model properly low mass halos orbiting within hosts of Milky Way mass.  This situation has recently improved with simulations that are now able to successfully sample the mass function on these scales, such as the Via Lactea, CLUES, GHALO and Aquarius simulations \\citep{Diemand2007, Diemand2008, Springel2008, Stadel2009, Libeskind2010}. For instance, using the Via Lactea simulations \\cite{Kuhlen2007} found that subhalos are also not spherical, although the effect of tides tends to make subhalos rounder than comparable objects in the field.  These results prompt questions about the validity of some of the assumptions involved in the mass modeling of stellar kinematics in dwarfs. And although the orbital anisotropy of the stars in a dSph is likely unrelated to that of dark matter (and associated to the formation history), it is nonetheless valuable to explore the dynamical structure of subhalos because they provide the main potential.  A detailed study of the shape of the Milky Way mass Aquarius halos was presented in \\citet{Vera2011}. Here we extend this analysis to lower mass objects, both subhalos of the main central halo and field halos, up to $1.5 h^{-1} \\rm Mpc$ from the center of the main Aquarius halo.  The paper is organized as follows. In Section \\ref{sec:numerical-modelling} we describe the numerical simulations, introduce the methods we use to determine halo shapes and explore the convergence of the results.  In Section \\ref{sec:shape-today} we compare the properties of subhalos and isolated objects of similar mass. We analyze subhalo shapes in the context of the kinematic modeling of dwarfs around the Milky Way in Section \\ref{sec:shape-mw} and summarize our main results in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   We have used the Aquarius simulations to study the shapes of field and satellite dark matter halos with emphasis on the mass range expected for the hosts of the dwarf galaxies in the Local Group. We have used an iterative method based on the normalized inertia tensor to characterize the principal axis lengths $a \\ge b \\ge c$ of halos and subhalos as a function of radius. In particular, we have explored in detail halo shapes measured in the inner regions (radius of maximum circular velocity $r_{\\rm max}$) and in the outskirts or \\emph{$r_{95}$ contour}. Although stars are more centrally concentrated than the dark matter, our resolution allows us to characterize halo shapes at radii as small as $r \\sim 1 \\; \\rm kpc$, starting to probe the regime traced by the outer stars in dwarf galaxies.  Through a comparison of objects in common between the different resolution levels of the Aquarius simulations, we have noticed that simple number of particles cuts do not guarantee convergence in the measured halo shapes, especially in the inner regions. We find that instead the \\emph{convergence radius} $r_{\\rm conv}$ (defined as the threshold $\\kappa = 7$ in the ratio between the local relaxation time and the dynamical time at the virial radius \\citep{Power2003}) provides a good estimate of the radius where the axis ratios are robustly determined (with an error $< 8$\\%).  We have found that the typical axis ratios of isolated halos in the Aquarius simulations decrease with increasing mass, or equivalently maximum halo circular velocity $V_{\\rm max}$, i.e.\\ low mass objects tend to be more spherical than Milky Way-like objects.  These trends are well approximated by a relation between the axis ratio measured at the $r_{95}$ radius $c/a|_{r_{95}}$ and $V_{\\rm max}$, i.e.  $c/a|_{r_{95}} \\sim -0.021 \\log V_{\\rm max}$, while this relation is slightly steeper if the axis ratio is measured at $r_{\\rm max}$ contour in which case $c/a|_{r_{\\rm max}} \\sim -0.032 \\log V_{\\rm max}$.  The differences in the shapes of field versus satellite halos are small and within the intrinsic scatter of the samples. Nonetheless, at a fixed $V_{\\rm max}$, subhalos tend to have larger axis ratios than isolated objects in the field.  The similarity between subhalos and field objects is also apparent in the lack of significant trends in the axis ratios with distance to the main host halo, $d$. We find, however, that the alignment of the ellipsoids varies with $d$: dark matter halos at close distances from the host center tend to be oriented preferentially with their major axis pointing radially. The signal disappears only for $d \\gtrsim 2.5$ $r_{\\rm vir}$, where the orientations are consistent with random.  We have also focused on the properties of subhalos likely to host analogs of the \\emph{classical satellites} of the Milky Way ($-13.2 \\leq M_V \\leq -8.6$), according to the semianalytic model of galaxy formation run on the Aquarius suite by \\citet{Starkenburg2012}. Our analysis indicates that these galaxies are hosted by mildly triaxial dark matter objects with minor-to-major axis ratios $c/a\\approx 0.60$ and intermediate-to-major $b/a\\approx 0.75$ in the first kiloparsec with a clear trend towards becoming axisymmetric in the outskirts. Their internal orbital structure shows evidence of being affected by tidal forces from their hosts (i.e. the main Aquarius halos), since the velocity anisotropy becomes tangential with radius, in clear contrast to what is found for isolated systems. We have also found that this orbital structure may be modeled in the axisymmetric context, where the velocity anisotropy $\\beta_z \\sim 0 $ along the minor axis, and declines with distance along the major axis. These results may be used to motivate more realistic models of the subhalos hosting satellite galaxies like those observed around the Milky Way."
    },
    "1402/1402.2301_arXiv.txt": {
        "abstract": "We detect a weak unidentified emission line at $E=(3.55-3.57)\\pm0.03$ keV in a stacked \\xmm spectrum of 73 galaxy clusters spanning a redshift range $0.01-0.35$. MOS and PN observations independently show the presence of the line at consistent energies. When the full sample is divided into three subsamples (Perseus, Centaurus+Ophiuchus+Coma, and all others), the line is seen at $>3\\sigma$ statistical significance in all three independent MOS spectra and the PN ``all others'' spectrum. The line is also detected at the same energy in the \\chandra\\ ACIS-S and ACIS-I spectra of the Perseus cluster, with a flux consistent with \\xmm\\ (however, it is not seen in the ACIS-I spectrum of Virgo). The line is present even if we allow maximum freedom for all the known thermal emission lines. However, it is very weak (with an equivalent width in the full sample of only $\\sim 1$ eV) and located within 50--110 eV of several known faint lines; the detection is at the limit of the current instrument capabilities and subject to significant modeling uncertainties. On the origin of this line, we argue that there should be no atomic transitions in thermal plasma at this energy. An intriguing possibility is the decay of sterile neutrino, a long-sought dark matter particle candidate. Assuming that all dark matter is in sterile neutrinos with $m_s=2E=7.1$ keV, our detection in the full sample corresponds to a neutrino decay mixing angle $\\sin^2(2\\theta)\\approx 7\\times 10^{-11}$, below the previous upper limits. However, based on the cluster masses and distances, the line in Perseus is much brighter than expected in this model, significantly deviating from other subsamples. This appears to be because of an anomalously bright line at $E=3.62$ keV in Perseus, which could be an Ar \\textsc{xvii} dielectronic recombination line, although its emissivity would have to be 30 times the expected value and physically difficult to understand. In principle, such an anomaly might explain our line detection in other subsamples as well, though it would stretch the line energy uncertainties. Another alternative is the above anomaly in the Ar line combined with the nearby 3.51 keV K line also exceeding expectation by a factor 10--20. Confirmation with \\chandra and \\suzaku, and eventually \\astroh, are required to determine the nature of this new line. (APJ HAS THE ABRIDGED ABSTRACT) ",
        "introduction": "Galaxy clusters are the largest aggregations of hot intergalactic gas and dark matter. The gas is enriched with heavy elements (\\citet{mitchell76,serlemitsos77} and later works) that escape from galaxies and accumulate in the intracluster/intergalactic medium (ICM) over billions of years of galactic and stellar evolution. The presence of various heavy ions is seen from their emission lines in the cluster X-ray spectra. Data from large effective area telescopes with spectroscopic capabilities, such as \\asca, \\chandra, \\xmm\\ and \\suzaku, uncovered the presence of many elements in the ICM, including O, Ne, Mg, Si, S, Ar, Ca, Fe, and Ni \\citep[for a review see, e.g.,][]{bohringer2010}. Recently, weak emission lines of low-abundance Cr and Mn were discovered \\citep{werner2006,tamura2009}. Relative abundances of various elements contain valuable information on the rate of supernovae of different types in galaxies \\citep[e.g.,][]{loewenstein2013} and illuminate the enrichment history of the ICM \\citep[e.g.,][]{bulbul2012b}. Line ratios of various ions can also provide diagnostics of the physical properties of the ICM and uncover the presence of multi-temperature gas, nonequilibrium ionization states, and nonthermal emission processes such as charge exchange \\citep[CX,][]{paerels03}.  As for dark matter, 80 yr from its discovery by \\citep{zwicky1933,zwicky1937}, its nature is still unknown (though now we do know for sure it exists, from X-ray and gravitational-lensing observations of the Bullet Cluster; \\citet{clowe06}, and we know accurately its cosmological abundance, e.g., \\citet{hinshaw13}). Among the various plausible dark matter candidates, one that has motivated our present work is the hypothetical sterile neutrino that is included in some extensions to the standard model of particle physics (\\citet{dodelson1994} and later works; for recent reviews see, e.g., \\citet{abazajian2007,boyarsky2009}). Sterile neutrinos should decay spontaneously with the rate \\begin{equation} \\Gamma_\\gamma(m_s,\\theta) = 1.38\\times 10^{-29}{\\rm\\ s}^{-1}\\ \\left(\\frac{\\sin^2 2\\theta}{10^{-7}}\\right) \\left(\\frac{m_s}{1\\rm\\ keV}\\right)^5, \\label{eq:rate} \\end{equation} where the particle mass $m_s$\\/ and the ``mixing angle'' $\\theta$\\/ are unknown but tied to each other in any particular neutrino production model \\citep{pal1982}. The decay of sterile neutrino should produce a photon of $E=m_s/2$ and an active neutrino. The mass of the sterile neutrino may lie in the keV range, which would place its decay line in the range accessible to X-ray observations of dark matter aggregations, such as clusters of galaxies, nearby galaxies, and the Milky Way \\citep{abazajian2001a,abazajian2001b}. So far, searches in various types of massive systems have resulted only in upper limits \\citep[e.g.,][]{boyarsky2012, abazajian2012}.  Current X-ray archives of \\xmm, \\chandra\\ and \\suzaku\\ contain vast collections of galaxy cluster observations. Mining these databases can result in significant improvement in sensitivity to faint spectral features compared to individual cluster observations (as proposed, e.g., by \\citet{abazajian2001b}), with respect to both the statistical and (in a less obvious way) systematic or instrumental uncertainties. In this paper, we undertake a fishing expedition that combines the spectra of many bright clusters from the \\xmm\\ archive in order to search for any kind of faint, unidentified X-ray emission lines --- be they thermal lines from previously undetected rare elements in the ICM or the elusive sterile neutrino decay line.  To improve the sensitivity to weak spectral lines and reduce systematic effects, we stack the X-ray spectra from clusters at different redshifts in their rest frame, rescaling the photon energies to $z=0$. After blueshifting each cluster spectrum to $z=0$, any background lines or instrumental response artifacts are smeared out (since they occur in the detector frame), but any weak line intrinsic to clusters would be amplified and may become detectable in the combined spectrum. In this paper, we use this method to detect a previously unknown, extremely faint emission line at $E\\approx 3.55-3.57$ keV. The line was detected in the stacked \\xmm X-ray spectra of 73 bright galaxy clusters in the redshift range $0.01< z< 0.35$, and independently in several subsamples. The redshift range in our sample corresponds to an energy difference between 0 and 1.2 keV at the energy of the line of $E\\approx 3.5$ keV. In the blueshifted frame, the response features, background, and instrumental lines are smeared by this line shift, while the detected line appears at the same energy, making it possible to detect the 3.57 keV line with high confidence.   This paper is organized as follows. Sections 2 and 3 describe the \\xmm data processing, background modeling, and spectra stacking methods. We also determine the best-fit flux and energy of the detected spectral feature using the \\xmm MOS and PN instruments. In Section 4, we confirm the existence of this spectral line in the \\chandra ACIS-I and ACIS-S spectra of the Perseus cluster, and obtain an upper limit from the ACIS-I observations of the Virgo Cluster. In Section 5, we discuss the possible interpretations of this new emission line. All errors quoted throughout the paper correspond to 68\\% (while in parentheses, 90\\%) single-parameter confidence intervals; upper limits are at 90\\% confidence, unless otherwise stated. ",
        "conclusions": ""
    },
    "1402/1402.6714_arXiv.txt": {
        "abstract": "We present a comprehensive high spatial-resolution imaging study of globular clusters (GCs) in NGC\\,1399, the central giant elliptical cD galaxy in the Fornax galaxy cluster, conducted with the {\\it Advanced Camera for Surveys} (ACS) aboard the {\\it Hubble Space Telescope} (HST).~Using a novel technique to construct drizzled PSF libraries for HST/ACS data, we accurately determine the fidelity of GC structural parameter measurements from detailed artificial star cluster experiments and show the superior robustness of the GC half-light radius, $r_h$, compared with other GC structural parameters, such as King core and tidal radius. The measurement of $r_h$ for the major fraction of the NGC\\,1399 GC system reveals a trend of increasing $r_h$ versus galactocentric distance, $R_{\\rm gal}$, out to about 10\\,kpc and a flat relation beyond. This trend is very similar for blue and red GCs which are found to have a mean size ratio of $r_{\\rm h, red}/r_{\\rm h, blue}\\!=\\!0.82\\pm0.11$ at all galactocentric radii from the core regions of the galaxy out to $\\sim\\!40$\\,kpc.~This suggests that the size difference between blue and red GCs is due to {\\it internal} mechanisms related to the evolution of their constituent stellar populations.~Modeling the mass density profile of NGC\\,1399 shows that additional {\\it external} dynamical mechanisms are required to limit the GC size in the galaxy halo regions to $r_h\\!\\approx\\!2$ pc. We suggest that this may be realized by an exotic GC orbit distribution function, an extended dark matter halo, and/or tidal stress induced by the increased stochasticity in the dwarf halo substructure at larger galactocentric distances. We compare our results with the GC $r_h$ distribution functions in various galaxies and find that the fraction of extended GCs with $r_h\\geq5$\\,pc is systematically larger in late-type galaxies compared with GC systems in early-type galaxies.~This is likely due to the dynamically more violent evolution of early-type galaxies.~We match our GC $r_h$ measurements with radial velocity data from the literature and split the resulting sample at the median $r_h$ value into compact and extended GCs. We find that compact GCs show a significantly smaller line-of-sight velocity dispersion, $\\langle\\sigma_{\\rm cmp}\\rangle\\!=\\!225\\!\\pm\\!25$ km s$^{-1}$, than their extended counterparts, $\\langle\\sigma_{\\rm ext}\\rangle\\!=\\!317\\!\\pm\\!21$ km s$^{-1}$.~Considering the weaker statistical correlation in the GC $r_h$-color and the GC $r_h$-$R_{\\rm gal}$ relations, the more significant GC size-dynamics relation appears to be astrophysically more relevant and hints at the dominant influence of the GC orbit distribution function on the evolution of GC structural parameters. ",
        "introduction": "\\subsection{Structural Parameters of Extragalactic GCs} Wide-field studies of massive galaxies provide important benchmarks for comparisons with glo\\-bu\\-lar cluster (GC) formation and evolution models as well as GC system assembly in the context of galaxy formation scenarios, not only because they define homogeneous and uniform datasets but also due to their simultaneous sampling of galaxy core and halo regions where various different physical processes affect the GC formation and survivability.~In general, GC formation is influenced by small-scale physics that governs star-formation and feedback processes \\citep[e.g.][]{murray92, harris94, ee97, hartwick09, murray09} while stellar feedback as well as internal and external dynamical mechanisms determine their early evolution \\citep{gieles08, bastianetal08, fall09, chandar09, elmegreen10, mapelli13} and the latter, ultimately, their fate \\citep[e.g.][]{gnedin97, vesperini97, vesperini03, chandar10, bekki10}.~The vast dynamical parameter ranges that need to be probed to study the complex interplay of these processes with numerical simulations are still very challenging for today's computers \\citep[e.g.][]{kravtsov05, li05, bournaud08, griffen10, schulman12, greif12}.~One simple approach to understand the influence of some of these processes on GC formation and evolution is the empirical study of GC structural parameters and their variation as a function of galactocentric distance.  Detailed GC structural parameters, such as core, half-light, and tidal radius, as well as central surface brightness, concentration, ellipticity, etc.~were, until the past decade, only accessible within the Local Group (LG) due to the limited spatial resolution of ground-based instrumentation \\citep[e.g.][]{king68, ill76, kontizas82, elson85, elson88, elson91, elson92, crampton85, demers90, trager95}.~The launch of the {\\it Hubble Space Telescope} (HST) catapulted this field to a whole new stratum, making vast numbers of GC systems accessible to high spatial resolution studies \\citep{harris13}. In fact, HST still provides our only access to high spatial resolution observations at optical wavelengths.~Several pioneering HST works quickly reached out with their GC half-light radius measurements beyond the LG as far as the Fornax galaxy cluster at $\\sim\\!20$ Mpc distance \\citep[e.g.][]{elson94, fusipecci94, kundu98, kundu99, puzia99, puzia00, zepf99}.~Numerous subsequent studies have used the superior spatial resolution of HST and the relatively large field of view ($\\sim\\!202\\arcsec \\!\\times\\! 202\\arcsec$) of the {\\it Advanced Camera for Surveys} (ACS) to collect large imaging datasets of extragalactic GC systems, the most homogeneous of which was obtained by the ACS Virgo and Fornax Cluster Surveys (ACSVCS and ACSFCS, see \\citealt{cote04} and \\citealt{jordan07}, respectively).~These observations set the baseline for systematic studies of GC structural parameters in the central regions of early-type cluster galaxies.~There was quickly mounting consensus among the early HST investigations that the observed central GCs had a rather broad half-light radius distribution with a peak somewhere in the range $\\sim\\!2\\!-\\!3$ pc, which led to the suggestion that this peak value may be used as a geometric distance indicator \\citep[e.g.][]{kundu01, jordan05}.~Another important finding was that the blue GCs are on average larger than the red GCs.~In particular, within the central regions of galaxies typically observed with HST, blue GCs show $\\sim\\!20\\%$ larger mean half-light radii compared to the red GC sub-population \\citep[e.g.][]{kundu98, kundu01, kundu99, puzia99, puzia00, zepf99, larsen01, jordan05, spitler06, harris06, harris09, harris10, blom12, goudfrooij12}.  However, one major limitation of most previous HST studies, targeting extragalactic GC systems such as the ACS Virgo and Fornax cluster surveys, was their limited field of view, using only one HST/ACS pointing per galaxy.~Because of this, these HST studies focused on the core regions of elliptical galaxies covering the inner few kpc (i.e.~$\\lesssim 1 R_{\\rm eff}$).~The outer parts of rich GC systems in central cluster galaxies were so far missed and mainly observed with ground-based instrumentation at much lower spatial resolutions \\citep[e.g.][]{rhode01, rhode04, rhode07}.~The only other ground-based study featuring a wide field of view {\\it and} high spatial-resolution was performed by \\cite{gomez07} using Magellan/IMACS under exceptional $\\sim\\!0.5\\arcsec$ average seeing conditions to measure half-light radii of 364 radial-velocity confirmed GCs in NGC~5128 (S0/E) out to $\\sim\\!8 R_{\\rm eff}$ of the spheroid light and found no significant correlation between GC half-light radius and projected galactocentric distance, i.e.~$r_h\\!\\propto\\!R^0$, outside $\\sim\\!1 R_{\\rm eff}$.~However, \\citeauthor{gomez07} reported that at $\\lesssim\\!1 R_{\\rm eff}$ the red GCs show a steeper $r_h\\!-\\!R$ relation and on average 30\\% smaller sizes than blue GCs.  Other studies using more than single-pointing HST observations conducted GC half-light radius measurements in NGC\\,4594 (Sa) out to about $6\\,R_{\\rm eff}$ of the bulge component \\citep[658 GC candidates]{spitler06, harris10}, in NGC\\,4365 (E) out to $\\sim\\!2.4 R_{\\rm eff}$ of the spheroid light \\citep[659 GC candidates]{blom12}, and in 6 giant elliptical galaxies out to $\\sim\\!4\\!-\\!5 R_{\\rm eff}$ of their spheroids \\citep[altogether 3330 GC candidates]{harris09}.~In the case of NGC\\,4594, the inner red GCs are $\\sim\\!17\\%$ smaller than the blue ones, but because of a steeper size-radius relation of the red GC sub-population, this difference becomes insignificant at galactocentric radii $\\gtrsim\\!2.7 R_{\\rm eff}$ of the bulge light.~NGC\\,4365 hosts on average $\\sim\\!32\\%$ larger blue GCs compared to their red counterparts and shows a steep size-radius relation, $r_h\\!\\propto\\!R^{(0.49\\pm0.04)}$, for the entire GC sample, similar to Milky Way's GC system \\citep[][]{vdB91}.~However, \\citeauthor{blom12}~do not investigate whether this relation differs between GC sub-populations as a function of projected galactocentric radius.~The composite GC system of the six giant ellipticals studied by \\citeauthor{harris09} exhibits a mild relation of the form $r_h\\!\\propto\\!R^{0.11}$ and a $\\sim\\!17\\%$ size difference between red and blue GCs that is independent of projected galactocentric radius.  \\subsection{Astrophysical Implications} In general, the finding of a size difference between blue and red GCs has important astrophysical implications for the understanding of the formation and evolution of GCs and for the usefulness of the peak value of the GC size distribution as geometric distance indicator. Several studies, such as \\cite{larsen03}, \\cite{jordan04}, and \\cite{harris09} put forward models to explain the size difference between blue and red GCs.~Inspired by the Milky Way GC system where a shallow relation exists between GC half-light radius and the 3-dimensional galactocentric distance, $r_h\\!\\propto\\!R_{\\rm 3D}^{0.5}$ \\citep{vdB91}, \\citeauthor{larsen03} suggested that the GC size difference between red and blue GCs in massive ellipticals could be due to the difference in their spatial distribution functions.~Typically, the red GC sub-population would be more centrally concentrated than their blue counterpart and therefore on average smaller, being tidally more truncated by the stronger host galaxy potential.~However, \\cite{webb12b} have shown in detailed numerical simulations that the observed GC size difference is unlikely due to projection effects alone.~In contrast to this {\\it external effect}, two alternative {\\it internal effects} were put forward. Firstly, \\cite{jordan04} suggested that the combined effect of mass segregation and shorter stellar lifetimes of more metal-rich stars at a given mass may explain the GC size difference.~This was strictly valid under the assumption that the GC {\\it half-mass} radius distribution would be independent of metallicity and that metal-poor and metal-rich GCs were of the same age, which may be at odds with observations \\citep[e.g.][]{puzia02, marinfranch09, goudfrooij12}.~This scenario was further developed by \\cite{sippel12} and \\cite{schulman12} in direct-integration N-body simulations of young, low-mass clusters with and without initial mass segregation, the absence of which was found to be enhancing the GC size difference.~\\cite{downing12} performed Monte-Carlo N-body simulations of massive star clusters and found that significant numbers of massive stellar remnants, i.e.~single and binary black holes would boost this GC size difference.~Secondly, \\cite{harris09} suggested that more metal-rich proto-GC clouds could cool more efficiently and therefore collapse into a more concentrated quasi-equilibrium state before forming stars than clusters formed from low-metallicity gas.~Any of these three scenarios comes with limiting assumptions and is likely not the single cause for the measured GC size difference as the variety of results described above indicates.  To provide a larger and statistically robust dataset to constrain GC sizes as a function of galactocentric radius, we embarked on a wide-field observing campaign covering a large area with an HST/ACS mosaic out to several effective radii ($>\\!5 R_{\\rm eff}$) of the diffuse spheroid light around NGC\\,1399, the central galaxy in the Fornax galaxy cluster that hosts one of the richest ($\\ga6000$ GCs; Specific frequency\\footnote{The specific frequency of a GC system is defined as twice the number of GCs brighter than the turn-over luminosity of the GC luminosity function, given by $N_{\\rm GC}$, relative to the absolute $V$-band luminosity of the host galaxy, $M_V$, which is normalized to $-15$ mag.~This quantity is defined as the specific frequency of a GC system $S_N\\!=\\!N_{\\rm GC}10^{0.4(M_V+15)}$; see also \\cite{georgiev10} and \\cite{harris13} for other GC system scaling relations.} $S_N\\approx5$) and most extended GC systems in the nearby Universe \\citep{dirsch03, faifer04, bassino06}. A significant part of the outer-halo GC system of NGC\\,1399 is located hundreds of kpc away from its host and is probing the transition regime between galaxy and cluster potential \\citep{ferguson89}.~At the same time, the formation efficiencies of these outer-halo, blue GCs appear to be higher than those of the inner red GCs, $S_N(\\mbox{red})\\!\\approx\\!3$ while $S_N(\\mbox{blue})\\!\\approx\\!14$ \\citep{forte05}. Spectroscopic radial-velocity studies of hundreds of GCs established a very complex multi-component system with the blue GCs being kinematically distinct from the red GC sub-population the latter of which shows dynamics similar to that of the host galaxy diffuse stellar component.~The blue GCs, on the other hand, seem to have been partly accreted from satellite galaxies \\citep{schuberth10}.~It is this large auxiliary kinematic dataset that makes the GC system of NGC\\,1399 an ideal target for a wide-field, high spatial-resolution study with HST/ACS \\citep[in comparison to M87, e.g.][]{epeng09, madrid09} as several hundreds of member stellar systems are robustly separated from the fore- and background in radial velocity space.  In our previous works, we used the dataset from this paper to study the Low Mass X-ray Binary (LMXB) population and the correlation of their properties with GC structural parameters \\citep{paolillo11, dago13}, as well as the GC selection techniques based on neural algorithms \\citep{brescia12}. Here we focus on the properties of the GC system itself.~Our present paper is organized as follows: in \\S2 we present the HST/ACS observations and discuss the details of sub-pixel dithering, \\S3 includes a description of the preliminary photometry that enters our structural parameter fitting code, which is introduced and thoroughly tested in \\S4.~We present our results in \\S5, where we show the large-scale variations of GC structural parameters within NGC\\,1399.~We discuss the implications in \\S6 and conclude this work in \\S7. ",
        "conclusions": "\\label{ln:discussion}  \\subsection{The Inner vs.~Outer GC System of NGC\\,1399} \\label{txt:innerouter} The significant GC size-luminosity relation of the inner $10$\\,kpc which disappears in the outer regions may indicate a transition in the predominance of various mechanisms at different galactocentric radii that shape the GC sizes and thus their evolution as a system.~Since the transition does not depend on GC color, i.e.~blue and red GCs show that same $r_h-R_{\\rm gal}$ relation, external dynamical effects are the most probable explanation (e.g.~dynamical friction of massive GCs that quickly sink into the core regions of the inner galaxy, tidal harassment of low-mass GCs by dwarf haloes in the outer halo regions, etc.).~While detailed numerical modelling of these effects goes beyond this work, we point out that our dataset is ideal to conduct detailed analyses such as those presented in \\cite{vesperini03} and \\cite{webb12a,webb13}.~We note that the mean $r_h$ for all resolved sources within $20\\arcsec\\!<\\!R_{\\rm gal}\\!<\\!120$\\arcsec\\ is $1.95\\pm0.06$ pc, i.e.~significantly smaller than the mean value for the entire GC system of $\\langle r_h\\rangle\\!=\\!3.21\\pm0.07$ pc.  \\begin{deluxetable}{lccc}[!t] \\centering \\tabletypesize{\\scriptsize} \\tablecaption{Photometric selection of blue and red GCs.\\label{tab:col_sel}} \\tablehead{\\colhead{} & \\colhead{blue GCs} & \\colhead{red GCs} & \\colhead{Ref.}} \\startdata Ground-based & $T1<23$& $T1<23$ &  (1) \\\\ data & $1.0\\leq C\\!-\\!R<1.65$ & $1.65\\leq C\\!-\\!R<2.2$ &\\\\\\\\ HST data & $z<22.5$ & $z<22.5$ & (2) \\\\ & $1.3\\leq g\\!-\\!z<1.9$ & $1.9\\leq g\\!-\\!z<2.5$ & \\\\ \\enddata \\tablerefs{(1): \\cite{bassino06}, (2): \\cite{kundu08}} \\end{deluxetable}  To test whether the stellar mass distribution in NGC\\,1399 is sufficient to produce the GC $r_h\\!-\\!R_{\\rm gal}$ trend (see Figure~\\ref{pdf:rh_galrad}) we use the surface brightness profile data obtained as part of the Carnegie-Irvine Galaxy Survey \\citep[CGS, see][]{ho11,li11} to compute the local {\\it instantaneous} Jacobi radius of GC ($r_J$) as a function of galactocentric radius out to $\\sim\\!280\\arcsec$ (i.e.~$\\sim\\!28$\\,kpc) which corresponds to the maximum sampling radius of CGS. The Jacobi radius marks the point at which the gravitation forces exerted on GC member stars due to the GC potential and that of its host galaxy are equal but opposite in direction. The Jacobi radius can be expressed as \\begin{equation} r_J=R_{\\rm gal}\\left(\\frac{m_{\\rm GC}}{2{\\cal M}_{\\rm gal}}\\right)^{1/3} \\label{eq:jacobi} \\end{equation} and is a robust representation of the {\\it instantaneous} GC tidal radius that is induced by the surrounding tidal field \\citep{innanen83,bertin08,renaud11,webb13}.  We proceed with computing the NGC\\,1399 mass distribution profile using the CGS data\\footnote{http://cgs.obs.carnegiescience.edu/CGS/Home.html} and the recipes outlined in \\cite{bell03} to convert photometric colors into stellar mass-to-light ratios as a function of galactocentric radius \\citep[see also][]{zibetti09,into13}.~We compute the ${\\cal M}_\\star/L_V$ profile through linear interpolation of predictions for a 13 Gyr old stellar population with variable metallicity, which is set by the measured photometric color profile of NGC\\,1399, using the 2007 update of the \\cite{bc03} SSP models. With the radial trend for ${\\cal M}_\\star/L_V$ we derive then the corresponding relation for the stellar mass \\begin{equation} \\label{eq:massprof} \\log({\\cal M}_\\star)\\!=\\!\\log({\\cal M}_\\star/L_V)_r\\!-\\!0.4[m_V(r)\\!-\\!D_L\\!-\\!M_{V,\\odot}] \\end{equation} enclosed in $R_{\\rm gal}\\!=\\!r$, where $m_V(r)$ is the integrated magnitude derived from the galaxy surface brightness profile, $D_L$ is the luminosity distance, and $M_{V,\\odot}\\!=\\!4.83$ mag is the absolute $V$-band magnitude of the Sun. With Equation~\\ref{eq:jacobi} and the derived stellar mass profile of NGC\\,1399 from Equation~\\ref{eq:massprof}, we determine the {\\it instantaneous} Jacobi radii for GCs with a total mass of $m_{\\rm GC}\\!=\\!10^3, 10^4$, and $10^5 M_\\odot$ using the results from \\cite{baumgardt10} and \\cite{ernst13} who determine the typical ratios between half-light and Jacobi radius with minimum, mean, and maximum values of $\\log(r_h/r_J)\\!=\\!-1.5,-1.0$, and $-0.5$ for Milky Way GCs. The extremes of the $r_h/r_J$ distribution are representative of GCs that are under- and overfilling their Roche lobes, respectively.  We also compute the GC stellar masses using the differential ${\\cal M}_\\star/L$ predictions from the GALEV SSP models \\citep{kotulla09}, assuming uniformly old GC ages ($t_{\\rm GC}=13$\\,Gyr) and using the $g\\!-\\!z$ and $C\\!-\\!T1$ GC colors to account for ${\\cal M}_\\star/L$ variations as a function metallicity. For GCs which lack color information we adopt the median ${\\cal M}_\\star/L$ of the GC sample for which photometric colors are available.  We overplot the corresponding expectation trends for $r_h$ as a function of galactocentric radius in Figure~\\ref{pdf:rh_galrad_jacobi} and use the color shading to parametrize GC mass.~Consistent with Figure~\\ref{pdf:rh_intmag} we see no preferred GC mass scale at a given galactocentric radius. We observe that none of the curves reproduces the break at 10\\,kpc of the $r_h\\!-\\!R_{\\rm gal}$ profile and its flatness at large galactocentric radii. The stellar mass density distribution is clearly not sufficient and requires an additional mechanism to limit GC sizes at large $R_{\\rm gal}$.~This could in principle be achieved by an exotic eccentricity distribution function of GC orbits, which would bring the outer clusters into the inner galaxy on preferentially radial orbits \\citep[see][]{webb13}.~An alternative explanation for the observed situation could be an additional tidal limitation of GCs in the outskirts of the galaxy, which could be realized in two different ways:  \\noindent 1) by an additional mass component in the form of a dark matter density profile of the NFW type $\\rho(r)\\!=\\!\\rho_0/[(r/R_s)(1+r/R_s)^2]$ \\citep{navarro96} where the total mass inside radius $R_{\\rm gal}$ is given by \\begin{eqnarray} {\\cal M}_{\\rm DM}(<\\!R_{\\rm gal})&=& 4\\pi \\int_0^{R_{\\rm gal}}r^2 \\rho (r) dr\\\\ &=&4\\pi \\rho_0 R_s^3 \\left[\\ln\\left(\\frac{R_s+R_{\\rm gal}}{R_s}\\right)-\\frac{R_{\\rm gal}}{R_s+R_{\\rm gal}}\\right]\\!. \\nonumber \\end{eqnarray} The resulting relations for $\\rho_0\\!\\approx\\!4\\cdot10^7 M_\\odot/{\\rm pc}^3$, $R_s\\!\\approx\\!130$\\,kpc, and $\\log(r_h/r_J)\\!=\\!-1.0$ are illustrated in Figure~\\ref{pdf:rh_galrad_jacobi} as magenta curves and show that even in the presence of a typical dark matter halo, i.e.~considering ${\\cal M}_\\star\\!+\\!{\\cal M}_{\\rm DM}\\!=\\!{\\cal M}_{\\rm gal}$ in Equation~\\ref{eq:jacobi}, the $r_h\\!-\\!R_{\\rm gal}$ GC relations are still monotonically increasing, albeit not as rapidly as in the case of considering ${\\cal M}_\\star$ only. Hence, an additional component is required to flatten out the $r_h\\!-\\!R_{\\rm gal}$ profiles at large galactocentric radii.  \\noindent 2) We, therefore, suggest that an increased stochastic distribution of low-mass dark matter haloes that are part of the galaxy cluster potential induce additional tidal stress on outer-halo GCs.~Such a changing mass fraction in subhaloes as a function of galactocentric radius is observed in high-resolution $\\Lambda$CDM simulations \\citep[e.g.][]{springel08} and would increase the ``tidal variance\" in outer-halo regions, thereby truncating the GC stellar density profiles.~This may limit the GC sizes to a roughly constant value, something that shall be explored with dedicated high-resolution numerical simulations.  \\subsection{Structural Parameter Distributions} \\label{ln:distribs}  \\begin{deluxetable*}{lccccccccccccc}[!t] \\centering \\tabletypesize{\\scriptsize} \\tablecaption{Fractions of extended GCs for various GC systems.\\label{tab:radfrac}} \\tablehead{\\colhead{Host Galaxy} & \\colhead{$E_5$} & \\colhead{$\\hat{E}_5$} &\\colhead{Ref.} &\\colhead{Dist./Mpc} &\\colhead{Ref.} & \\colhead{$R_{\\rm gal}$/kpc} & \\colhead{$E_{5/10}$} & \\colhead{$\\hat{E}_{5/10}$} &\\colhead{$r_e(K_s)$} &\\colhead{$r_e$/kpc} & \\colhead{${\\cal E}_{5}$} & \\colhead{$\\hat{{\\cal E}}_{5}$} } \\startdata NGC~1399 \t\t\t& 0.122 & 0.62 & (1) & $20.13\\pm0.4$ & (7) & 51.3 & 0.061 & 0.21 & 32.9\\arcsec & 3.21& 0.0480 & 0.12 \\\\ NGC~4486 (M87)\t\t& 0.066 & 0.34 & (2) & $16.70\\pm0.2$ & (8) & 12.3 & 0.064 & 0.22 & 41.5\\arcsec & 3.36& 0.0633 & 0.16 \\\\ NGC~4472 (M49)\t\t& 0.073 & 0.37 & (2) & $16.40\\pm0.2$   & (8) & 11.6 & 0.072 & 0.25 & 56.1\\arcsec & 4.46& 0.0714 & 0.18 \\\\ NGC~4594 (M104)\t\t& 0.026 & 0.13 & (3) & $9.08\\pm0.2$   & (9) & 15    & 0.024 & 0.08 & 55.3\\arcsec & 2.43& 0.0160 & 0.04 \\\\ NGC~5128 (Cen\\,A)\t\t& 0.170 & 0.86 & (4) & $3.84\\pm0.35$ &(10)& 20    & 0.179 & 0.61 & 82.6\\arcsec & 1.54& 0.2127 & 0.54 \\\\ NGC~224   (M31)\t\t& 0.241 & 1.22 & (5) & $0.779\\pm0.05$&(11)& 160  & 0.132 & 0.45 & 443.2\\arcsec & 1.67& 0.1316 & 0.33 \\\\ Milky Way            \t\t& 0.197 &$\\equiv1$&(6)& $\\dots$&$\\dots$&120&0.292&$\\equiv1$&$\\dots$&2.50&0.3974&$\\equiv1$\\\\ \\enddata \\tablecomments{$R_{\\rm gal}$ is the maximum sampling radius of the corresponding dataset in kpc.~$E_5$ and $\\hat{E}_5$ are the values defined in Equations~\\ref{eq:S5} and \\ref{eq:S5hat}, while the corresponding values for the GC samples restricted to $R_{\\rm gal}\\!\\leq\\!10$\\,kpc are given as $E_{5/10}$ and $\\hat{E}_{5/10}$ and those within 2.5 effective radii as ${\\cal E}_{5}$ and $\\hat{{\\cal E}}_{5}$, respectively.~$K_s$-band effective radius measurements, $r_e(K_s)$, are from 2MASS and were obtained from the NASA/IPAC Infrared Science Archive.~For the Milky Way, the corresponding value was adopted based on the predictions of the Besan\\c{c}on Galactic stellar population synthesis model \\citep{robin03}.} \\tablerefs{For the GC populations, (1): this work, (2): ACSVCS, see \\cite{jordan09}, (3): \\cite{harris10}, (4): \\cite{woodley10}, (5): \\cite{peacock09} and \\cite{huxor14}, (6): McMaster catalog, 2010 update of \\cite{harris96}. For the distance measurements, (7): \\cite{dunn06}, (8): \\cite{mei07}, (9): \\cite{jensen03}, (10): \\cite{harrisG10}, (11): \\cite{conn12}.} \\end{deluxetable*}  We show the $r_h$ distribution of NGC\\,1399 GCs in Figure~\\ref{pdf:h_distr} together with corresponding measurements for Milky Way and M31 GCs, taken from the McMaster catalog \\citep[2010 update of][]{harris96} as well as \\cite{peacock09} and \\cite{huxor14}, respectively.~In addition, we compare our half-light radius measurements to the $r_h$ distributions of GCs in NGC~5128 \\citep{woodley10}, the Sombrero galaxy \\citep[M104,][]{harris10}, and the two brightest Virgo ellipticals M49 and M87 which were studied by the ACSVCS \\citep[for details see][]{jordan09}.  The bottom panel of Figure~\\ref{pdf:h_distr} shows the entire sample of NGC\\,1399 GCs together with the distribution of radial-velocity confirmed GCs (dark histogram) and stars (cyan histogram).~It is important to note that all spectroscopically confirmed foreground stars concentrate around $r_h\\!\\approx\\!0$ pc, where unresolved objects are generally expected.~We provide mean and median values of each $r_h$ distribution in each panel of Figure~\\ref{pdf:h_distr} and point out that there is a trend of decreasing $r_h$ with increasing host galaxy luminosity \\citep{masters10} in which NGC\\,1399 and its central GC system fit right in. Such a trend generally supports the notion that the host environment has an impact on the GC $r_h\\!-\\!R_{\\rm gal}$ relation (see discussion above), and will depend on the sampled $R_{\\rm gal}$ range.  We also compare our sample to the measurements of \\cite{masters10} who derived GC half-light radii for the central regions in NGC\\,1399 from the ACS Fornax Cluster Survey data \\citep{jordan07} which, similar to the ACSVCS in Virgo, sampled massive early-type galaxies in Fornax with one central HST/ACS pointing. Both our and the ACSFCS distributions show very similar shapes and drop-offs from $\\sim\\!1.5$ pc up to about 5 pc, beyond which our sample starts to include many more extended GCs. This is mainly due to the nine times larger field of view of our data and we point out that many of these extended sources are radial-velocity confirmed {\\it bona-fide} GCs at large galactocentric radii.  The comparison with other GC systems in the upper panels of Figure~\\ref{pdf:h_distr} shows that all half-light radius distributions have very similar shapes featuring a relatively steep increase in GC number density at low $r_h$ values, with a peak somewhere in the range of $2\\!-\\!3$ pc, and a shallower decline towards more extended objects.~This distribution is dependent on the sampling of the GC luminosity function, galactocentric radius, as well as the amount of contamination, the measurement errors, and the GC selection criteria \\citep[e.g.][]{brescia12}.~It is hard to compare the unresolved parts at $r_h\\!\\la\\!1$ pc for galaxies further away than Sombrero (NGC\\,4594 at $D\\!\\approx\\!9$ Mpc) due to the resolution limit of HST (see shaded regions in Figure~\\ref{pdf:h_distr}).~Despite this limitation there is ample information and some intriguing aspects of the GC size distributions for sources with $r_h\\!\\ga\\!1.5$ pc.~In NGC\\,1399, these extended clusters predominantly reside at projected galactocentric radii, $R_{\\rm gal}$, larger than 10 kpc (see Figure~\\ref{pdf:rh_galrad}).~Since the observed GC populations in M49, M87, and M104 are all inside this radius (see Table~\\ref{tab:radfrac}), we find a very small population of similarly extended GCs in the corresponding samples.~This is, of course, an observational bias considering our and the earlier results by \\cite{vdB91} and \\cite{larsen03} who found correlations of the type $r_h\\!\\propto\\!R_{\\rm gal}^{n}$ with $n\\!<\\!1$, and \\cite{jordan05} who suggested an analytic expression that approximates the $r_h$ distribution for the inner GC systems in Virgo ellipticals.  Having sampled a significant population of GCs to large galactocentric radii in NGC\\,1399 in combination with similar results for less rich GC system (see Figure~\\ref{pdf:h_distr}), we, therefore, suggest that all GC systems are comprised of two components of clusters: one standard GC population with a size distribution resembling the typical GC half-light radius of $2\\!-\\!3$ pc and a second, less rich component of more extended GCs that are predominantly found at larger galactocentric radii. Alternatively, there might be a combination of mechanisms (explored further below) that act on just one GC population, but their effects manifest themselves at different radii, so that the extended GCs are only observed at large $R_{\\rm gal}$.  To quantify the fraction of extended GCs in a GC system, we define the number ratio of GCs with sizes larger than 5 pc relative to the total GC population, \\begin{equation} \\label{eq:S5} E_5=N_{\\rm GC}(r_h\\!\\geq\\!5 {\\rm pc})/N_{\\rm GC}({\\rm all}), \\end{equation} and normalize this value to the Galactic GC system, i.e. \\begin{equation} \\label{eq:S5hat} \\hat{E}_5=\\frac{ N_{\\rm GC}(r_h\\geq5 {\\rm pc})}{N_{\\rm GC}({\\rm all})} \\left(\\frac{N_{\\rm GC}(r_h\\geq5 {\\rm pc})_{\\rm MW}}{N_{\\rm GC} ({\\rm all})_{\\rm MW}}\\right)^{-1}. \\end{equation} The results for all GC systems are summarized in Table~\\ref{tab:radfrac} for the galactocentric sampling ranges of the corresponding dataset, which vary by about an order of magnitude.  In order to representatively compare the GC samples we, therefore, restrict each dataset to within $R_{\\rm gal}\\!\\leq\\!10$\\,kpc (about the maximum homogeneous sampling radius of the samples) as well as 2.5 effective radii of the host galaxy's diffuse light (set by the maximum radial sampling of each dataset), measured in the near-infrared $K_s$ filter.~We summarize the corresponding values as $E_{5/10}$ and $\\hat{E}_{5/10}$ for $R_{\\rm gal}\\!\\leq\\!10$\\,kpc as well as ${\\cal E}_{5}$ and $\\hat{{\\cal E}}_{5}$ for $R_{\\rm gal}\\!\\leq\\!2.5\\,r_e$ in Table~\\ref{tab:radfrac}.  We find a clear dichotomy in the ${\\cal E}_{5}$ (and $E_{5/10}$) between late-type and early-type galaxies.~While the three giant ellipticals NGC\\,1399, M87 and M49 as well as M104 show ${\\cal E}_{5}$ values clearly below 10\\%, the two late-type spirals, i.e.~M31 and the Milky Way, as well as NGC\\,5128 stand out with significantly higher ${\\cal E}_{5}$ values, clearly above $\\sim\\!10\\%$.~We attribute this result to differences in the tidal environment properties throughout the dynamical evolution and merging history of these galaxies.~Giant ellipticals experience in general a more violent evolution than spirals. It is unclear yet, how these numbers compare to other GC systems, but the fact that ${\\cal E}_{5}$ values of NGC\\,1399 and the two Virgo giant ellipticals M87 and M49 are remarkably similar, hints at physical processes acting that are acting in a similar way on the size evolution of their GC systems.~This includes the somewhat surprising result for the Sombrero galaxy's GC system with a similar ${\\cal E}_{5}$ value as the giant ellipticals. Higher ${\\cal E}_{5}$ values for the Milky Way and M31 might be the result of the dynamically more benign tidal field around such distant GCs and/or the younger, i.e.~less evolved, nature of NGC\\,5128, a recent merger remnant, and its GC system. How these numbers will play out for the GC systems in other Virgo cluster galaxies will be shown by the {\\it Next Generation Virgo Cluster Survey} (NGVS) which achieves a spatial resolution of $\\sim\\!5$\\,pc for the entire Virgo galaxy cluster out to its virial radius \\citep{ferrarese12, munoz14}. At least then it will be clear whether late-type galaxies have a systematically larger population of extended GCs than early-type galaxies, which host GC systems with a relatively smaller population of extended GCs.  Of course, we expect a complex interplay between the formation paths of the compact and extended GCs. In fact, we expect multiple components in the GC size distribution depending on the star cluster formation history and the evolution of the host galaxy. However, in a simplistic picture we speculate that while the primary component GCs (i.e.~compact GCs) are likely massive and old, and formed {\\it in-situ}, the nature of secondary component GCs (i.e.~extended GCs) is likely the result of a combination of populations of 1) dissolving star clusters triggered by recent formation of younger, low-mass GCs combined with increased tidal stress, e.g.~in central regions of galaxy clusters or merger remnants \\citep{gieles11, goudfrooij12}, 2) the accretion of more extended GCs from satellite galaxies which formed and survived in a more benign tidal environment \\citep[see also][]{georgiev09, dacosta09, smith13}, and/or 3) disrupting cores of stripped dwarf galaxy nuclei \\citep[e.g.][]{oh00, bekki03, pfeffer13}.~The corresponding detailed analysis of this scenario is the focus of a forthcoming paper.  \\subsection{Kinematic Properties of Compact and Extended GCs}  \\begin{figure}[!t] \\centering \\includegraphics[width=8.9cm]{fig17_rh_galrad_LOSVD.jpg} \\caption{({\\it Top panel}): The plot shows all measurements for individual GCs as in Figure~\\ref{pdf:rh_galrad} with the green dash-dotted curve approximating the running median relation as defined in Equation~\\ref{eqn:split}.~We use this empirical separation to define formally compact and extended GCs.~({\\it Middle panel}): The radial velocity of each GC matched with the \\cite{schuberth10} sample as a function of galactocentric radius.~The mean radial velocity of the sample is shown as horizontal dashed line.~The symbols are parametrized by the GC color (see Table~\\ref{tab:col_sel}) and split into blue and red GCs (shown in corresponding colors) and those without color information (shown as black dots). Green circles indicate GCs that have sizes formally more compact than relation shown as dash-dotted relation in the top panel.~({\\it Bottom panel}):~Sliding-median relations of the line-of-sight radial velocity dispersion as a function of galactocentric radius for compact (GC$_{\\rm cmp}$, {\\it green}) and extended GCs (GC$_{\\rm ext}$, {\\it black}) with their corresponding 90\\% confidence limits shown as dotted curves.} \\label{pdf:rh_galrad_losvd} \\end{figure}  The availability of matched GC size and radial velocity measurements in NGC\\,1399 allows us to investigate correlations between these two parameters. The radial velocities of all GCs with size measurements do not correlate in any statistically significant way as a function of projected galactocentric radius (Figure~\\ref{pdf:rh_galrad_losvd}).~We measure a total systemic heliocentric radial velocity of the entire sample as $\\langle v_{\\rm helio}\\rangle\\!=\\!1456\\!\\pm\\!17$ km s$^{-1}$ with a line-of-sight velocity dispersion of $\\sigma\\!=\\!295$ km s$^{-1}$ in good agreement with \\cite{schuberth10}.~This is also in good agreement and consistent with previous measurements of the diffuse light, i.e.~$\\langle v_{\\rm helio}\\rangle\\!=\\!1425\\!\\pm\\!4$ km s$^{-1}$ and $\\sigma_0\\!=\\!353\\pm19$ km s$^{-1}$ \\citep{graham98}, respectively.  Next, we divide our sample in a population of compact (GC$_{\\rm cmp}$) and extended GCs (GC$_{\\rm ext}$) using as division the relation illustrated as green dash-dotted line in the top panel of Figure~\\ref{pdf:rh_galrad_losvd}, which is approximating the running-median $r_h$ curve of the entire sample (black curves). This linear separation can be numerically expressed as \\begin{equation} r_h [{\\rm pc}] = \\left\\{ \\begin{array}{l l} 0.012\\,R_{\\rm gal} + 0.8 & \\quad \\text{if $R_{\\rm gal}<100$\\arcsec}\\\\ 2 & \\quad \\text{if $R_{\\rm gal}\\geq100$\\arcsec.} \\end{array} \\right. \\label{eqn:split} \\end{equation} We scrutinize the GC size-$v_{\\rm helio}$ relation for any correlations and find no significant slope for compact and extended GCs as a function of galactocentric radius $R_{\\rm gal}$.~However, looking at the line-of-sight velocity dispersion, $\\sigma$, of each of those sub-samples, we find a surprisingly clear dichotomy between compact and extended GCs in terms of their mean velocity dispersion.~While the compact GC sample exhibits $\\langle\\sigma_{\\rm cmp}\\rangle\\!=\\!225\\!\\pm\\!25$ km s$^{-1}$, we compute a much higher value for the extended sample with $\\langle\\sigma_{\\rm ext}\\rangle\\!=\\!317\\!\\pm\\!21$ km s$^{-1}$.~This is consistent with the $\\sigma$ differences found by \\cite{schuberth10} between the blue and red GC sub-population in NGC\\,1399 at a similar range in galactocentric radius.  Plotting the sliding median of the $\\sigma\\!-\\!R_{\\rm gal}$ relation (bottom panel of Figure~\\ref{pdf:rh_galrad_losvd}) reveals that this difference is most pronounced in the range $50\\arcsec\\!\\la\\! R_{\\rm gal}\\!\\la\\!320\\arcsec$, which roughly corresponds to the physical range of $15\\!\\la\\! R_{\\rm gal}\\!\\la\\!32$\\,kpc.~Outside this range, the difference seem to disappear, but we lack sample statistics to make definitive conclusions, and defer a more detailed analysis of this surprising result to a future study, when more comprehensive radial velocity samples become available.~Here we just note that given the scatter of the rather weak correlation between GC size and color (see Section~\\ref{ln:sizecol}), the significantly lower velocity dispersion of more compact (i.e.~red) GCs compared to their more extended counterparts (i.e.~blue GCs) appears to be the astrophysically stronger relation, which likely has its origin in the stronger influence of {\\it external} tidal truncation effects compared to {\\it internal} mechanisms that govern the GC size.~This is also consistent with our result of the flatter GC size-color relation for the inner vs.~outer GC sample (see Figure~\\ref{pdf:rh_color}).~These findings indicate the preferential influence of {\\it external} dynamical effects damping the size difference between red and blue GCs which is predominantly driven by {\\it internal} evolution of their constituent stellar populations and is likely a corollary of the GC size-dynamics correlation. Future GC radial velocity samples of the inner GC system in NGC\\,1399 will shed light on how the GC orbit distribution function influences these relations."
    },
    "1402/1402.3594_arXiv.txt": {
        "abstract": "{} % {We explore the properties of galaxies on the outskirts of clusters and their dependence on recent dynamical history in order to understand the real impact that the cluster core has on the evolution of galaxies.} {We analyse the properties of more than 1000 galaxies brighter than $M_{^{0.1}r}$=-19.6 on the outskirts of 90 clusters ($1 < r/r_{vir} < 2$) in the redshift range $0.05<z<0.10$. Using the line of sight velocity of galaxies relative to the cluster\u2019s mean, we selected high and low velocity subsamples. Theoretical predictions indicate that a significant fraction of the first subsample  should be backsplash galaxies, that is, objects that have already orbited near the cluster centre. A significant proportion of the sample of high relative velocity (HV) galaxies seems to be composed of infalling objects.} {Our results suggest that, at fixed stellar mass, late-type galaxies in the low-velocity (LV) sample are systematically older, redder and have formed fewer stars during the last 3 Gyrs than galaxies in the HV sample. This result is consistent with models that assume that the central regions of clusters are effective in quenching the star formation by means of processes such as ram pressure stripping or strangulation. At fixed stellar mass, LV galaxies show some evidence of having higher surface brightness and smaller size than HV galaxies. These results are consistent with the scenario where galaxies that have orbited the central regions of clusters are more likely to suffer tidal effects, producing loss of mass as well as a re-distribution of matter towards more compact configurations. Finally, we found a higher fraction of ET galaxies in the LV sample, supporting the idea that the central region of clusters of galaxies may contribute to the transformation of morphological types towards earlier types.} {} ",
        "introduction": "\\label{sec:intro}  The relative importance of nature vs. nurture (heredity vs. environment) in the formation and evolution of galaxies has been increasingly discussed in recent years. There is consensus that both are important and \\citet{DeLucia:2012} have shown that these two elements of galaxy evolution are inevitably intertwined. In the downsizing scenario, stars in more massive galaxies tend to have formed earlier and over a short period of time (\\citealt{Neistein:2006}). This effect in the early stages of galaxy formation suggests that the initial conditions (nature) are important.  The strong correlation between the present-day galaxy properties and the environment is well known (see for instance \\citealt{Blanton:2005}). In the hierarchical clustering scenario, the structures form from the bottom up, via the merging of small objects into successively more massive systems. In this scenario, the environment in which galaxies evolve can vary substantially with time. Extreme cases are clusters of galaxies, which favour a great number of physical processes that could modify the evolutionary history of galaxies. In this environment, gravitational phenomena such as the interactions and merger of galaxies are frequent and can produce morphological transformations \\citep{Moore:1998}. Moreover, the presence of the hot intra-cluster medium favours processes like ram pressure (\\citealt{GG:1972}, \\citealt{Abadi:1999}) or strangulation \\citep{Larson:1980}, which can affect both the gaseous content of galaxies and their star formation history \\citep{Fujita:1999}. Galaxies can also suffer the stripping of gas, stars and dark matter due to interaction with the cluster potential \\citep{Moore:1999}. These aspects make clusters of galaxies ideal objects for evaluating the processes involved in the evolution of galaxies. Recently it has been suggested that group environment also plays a key role in present-day galaxy properties (e.g. \\citealt{Wilman09}). \\citet{Poggianti09} analysed a combined sample of nearby and high redshift clusters of galaxies and found that the spiral and S0 fractions have evolved more strongly in less massive clusters, suggesting that morphological evolution since redshift 1 is more pronounced in low-mass clusters. \\citet{Rasmussen:2012} found that the fraction of star-forming group members is suppressed relative to the field, mirroring results for massive clusters (see also \\citealt{Mahajan:2013}).  The well-known morphological segregation of galaxies into clusters (\\citealt{Dressler:1980}) does not imply that the local density or distance to the cluster centre remains unchanged. Since the typical time-scale of galaxy orbits in a cluster is of the order of a few Gyrs (\\citealt{McCarthy:2008}), these galaxies must have passed close to the centre of the cluster since their infall. Moreover, it is to be expected that a fraction of the galaxies that have been close to the cluster centre, would today be on the outskirts. \\citet{Gill05} using high-resolution N-body simulations found that approximately half of the galaxies that at present have clustocentric distances between 1 and 2 virial radii, were in the past close to the cluster centre. These galaxies are usually called \"backsplash\". More recently, \\citet{Mahajan11} investigated how, at a given projected clustocentric radius, the stellar mass and star formation activity of a galaxy depend on its line-of-sight velocity relative to the cluster\u2019s mean velocity ($v_{los}$). Based on N-body simulations, they found that the fraction of backsplash objects depends on the velocity threshold: one-third (half) at projected radii slightly greater than the virial radius and $|v_{los}| < \\sigma$ ($|v_{los}| << \\sigma$), where $\\sigma$ is the velocity dispersion of the cluster. If the backsplash and the infalling populations can be identified separately, the effects on the galaxy properties caused by the core of clusters can be distinguished from the preprocessing in groups. \\citet{Mahajan11} analysed a sample of clusters selected from a group catalogue and found that star formation in galaxies is almost completely quenched in a single passage through the cluster. The role of the backsplash scenario in the evolution of galaxies is also tested in \\citet{Rines05}, \\citet{Pimbblet06} and \\citet{Aguerri10}.  Based on the backsplash scenario, in this paper we address the issue of the influence of cluster environment in galaxy evolution. For this purpose, we use the sample of galaxies in clusters selected by \\nocite{Coenda:2009} Coenda \\& Muriel (2009, hereafter paper I) plus a new subsample of X-ray clusters obtained by cross-correlating the Northern ROSAT All-Sky Galaxy Cluster Survey (NORAS, \\citet{Bohringer:2000}, hereafter B00) with the Sloan Digital Sky Survey (SDSS, \\citealt{York:2000}). We have analysed a wide range of properties of high and low velocity subsamples of cluster members with clustocentric distances $1<r/r_{vir}<2$ in the redshift range $0.05<z<0.10$. The results are analysed in the framework of models of the formation and evolution of galaxies.  This paper is organised as follows: in Sec. \\ref{sec:sample} we describe the samples of clusters and galaxies. Galaxy properties and their dependence on environment are analysed in Sec. \\ref{sec:results}. In Sec. \\ref{sec:conclu} we discuss the results and outline our conclusions. ",
        "conclusions": "\\label{sec:conclu}   Based on a sample of galaxy clusters composed of 90 X-ray selected and MaxBCG clusters in the redshift range $0.05<z<0.10$, we analyse properties of 1060 SDSS galaxies on the outskirt of clusters ($1 < r/r_{vir} < 2$). Using the radial velocity of galaxies (relative to the mean cluster velocity), we select low and high velocity subsamples as follows: $|v_{los}| < 0.5\\sigma$  and $|v_{los}| > \\sigma$ respectively. According to theoretical predictions, a significant fraction of the first subsample  should be backsplash galaxies, that is, objects that have already orbited near the cluster centre. In the backsplash scenario, the sample of galaxies with high relative velocity (HV) would include a significant fraction of objects that are falling into the cluster for the first time. We analyse two sets of properties. i) Photometric properties: the colour $^{0.1}(g-r)$, the Petrosian $r$-band absolute magnitude $M_{^{0.1}r}$, the radius that encloses 50\\% of the Petrosian flux and the $r$-band surface brightness computed inside $r_{50}$. ii) Properties derived from the spectral synthesis analysis: we use the VESPA algorithm (\\citealt{Tojeiro09}) to compute the stellar mass $M_{*}$, the mean age $\\tau$ weighted by mass, the metallicity $Z$ and the fraction of the stars that were formed during the last 3 Gyrs.    We find that ET-LV galaxies are systematically less massive than ET-HV, in agreement with the predictions of \\citet{Gill05} who found clear evidences that LV galaxies lose part of their mass in a single passage through the cluster centre. For the whole sample of galaxies, we find that the fraction of ETs grows with the stellar mass, going from 50\\% for the lowest mass bin and up to 80\\% for high mass galaxies. The fraction of ETs was larger in the LV sample than in the HV. The difference was larger for low mass galaxies, in which the fraction of ETs in the LV sample was up to 35\\% larger than in the HV sample. This result supports the scenario in which the central region of clusters of galaxies can contribute to the transformation of morphological types.  Galaxies with the same stellar mass and in the same environment are expected to be similar. Departures from the mean properties may be directly linked to differences in the recent dynamical history, as would be the case when comparing LV and HV galaxies with the same stellar mass in the same cluster region. We test this hypothesis for each of the galaxy properties considered in our work. Our results show that, for the same stellar mass, LV galaxies are on average older than HV.  If morphological types are considered, we find that this effect corresponds to LT galaxies. These results are consistent with the age gradients found in clusters by \\citet{Rakos:2007}. \\citet{Smith:2012} found that the age gradient depends on the mass of galaxy being lower for the most massive objects. These authors suggested that, even if there is a significant pre-processing in groups, the location in the cluster core at z = 0 effectively selects galaxies which passed through all stages of the environmental history, including pre-processing, earlier than those which today reside in the cluster outskirts (see also  \\citealt{Price:2011}). Within this scenario, LV are more advanced in the process of accretion and should therefore be, on average older, as observed in our analysis.  We find some evidences that LT-HV galaxies have formed more stars during the last 3 Gyrs than LT-LV galaxies, consistent with \\citet{Mahajan11} and with models that assume that a passage through the central region of clusters tends to quench star formation due to processes like ram pressure stripping or strangulation. This phenomenon is also reflected in the colours of galaxies. Our results show that LT-LV galaxies are less blue (have had less recent star formation) than LT-HV.  Although marginally significant, we find that LV galaxies with $\\log(M_{\\ast}/M_{\\odot})\\geqq 10.7$ tend to have a higher surface brightness than HV. Analysing the size and the luminosity of galaxies, our results suggest that the higher values of $\\mu_{50}$ for LV galaxies is likely to be due to smaller size of LV galaxies (for a given stellar mass) in comparison to HV galaxies. These results could be related to gravitational type processes, in which galaxies that have orbited the central regions of clusters are more likely to suffer tidal effects, such as the loss of part of their mass as well as a re-distribution of matter towards more compact configurations. This result is in agreement with \\citet{Aguilar86} who using N-body simulations found that strong collisions result in a reduced effective radius and a higher surface density."
    },
    "1402/1402.4461_arXiv.txt": {
        "abstract": "We use cosmological hydrodynamical simulations to investigate how the inclusion of physical processes relevant to galaxy formation (star formation, metal-line cooling, stellar winds, supernovae and feedback from Active Galactic Nuclei, AGN) change the properties of haloes, over four orders of magnitude in mass. We find that gas expulsion and the associated dark matter (DM) expansion induced by supernova-driven winds are important for haloes with masses $M_{200} \\la 10^{13}\\Msun$, lowering their masses by up to $20\\%$ relative to a DM-only model. AGN feedback, which is required to prevent overcooling, has a significant impact on halo masses all the way up to cluster scales ($M_{200} \\sim 10^{15}\\Msun$). Baryon physics changes the total mass profiles of haloes out to several times the virial radius, a modification that cannot be captured by a change in the halo concentration. The decrease in the total halo mass causes a decrease in the halo mass function of about $20\\%$. This effect can have important consequences for the abundance matching technique as well as for most semi-analytic models of galaxy formation. We provide analytic fitting formulae, derived from simulations that reproduce the observed baryon fractions, to correct halo masses and mass functions from DM-only simulations. The effect of baryon physics (AGN feedback in particular) on cluster number counts is about as large as changing the cosmology from WMAP7 to Planck, even when a moderately high mass limit of $ M_{500}\\approx10^{14}\\Msun$ is adopted. Thus, for precision cosmology the effects of baryons must be accounted for. ",
        "introduction": "\\label{sec:introduction} N-body simulations using only gravitationally interacting dark matter (DM) particles have played an important role in the development of astrophysical cosmology \\citep[e.g.][]{Frenk12}. DM simulations have for example been used to predict the large-scale distribution of matter, the halo mass function, and the density profiles of haloes. The results from such simulations form the basis for halo-based models and abundance matching techniques \\citep[e.g.][]{Seljak00,Cooray02}. DM simulations are also used as the starting point for semi-analytic models that associate galaxies to DM haloes in post-processing, and then follow the evolution of these galaxies according to different prescriptions that are calibrated such that the model reproduces a limited set of observables \\citep[e.g.][]{Baugh06}.  Neglecting the hydrodynamics and the feedback processes that affect the gas component also allows one to perform simulations with a dynamic range that would otherwise not be achievable due to the higher computational cost associated with the inclusion of baryon physics. Explicitly accounting for baryons in simulations means computing hydrodynamic forces and including processes like radiative cooling, photo-heating, star formation, metal enrichment, and also energetic feedback processes such as winds driven by supernovae and AGN that are able to generate outflows and eject baryons from (the inner parts of) DM haloes.  For many purposes it is a reasonable approximation to assume that the effect of baryon physics on the matter distribution is small, such as in the intergalactic medium (e.g., \\citealt{Theuns02,Viel12}) and on the outskirts of galaxy clusters where the gas has a long cooling time and approximately traces the dark matter (e.g., \\citealt{LeBrun14}).  However, it is clear that on small scales and in lower mass haloes, where cooling can allow the gas to condense to high densities, baryonic processes such as galactic winds can have important effects on the matter distribution.  Moreover, the distribution of the DM will itself adjust to the resulting change in the gravitational potential. Indeed, it appears that the observed rotation curves of dwarf galaxies cannot be reproduced by simulations that assume the standard cold dark matter paradigm unless they include the effect of outflows \\citep[e.g.][]{Governato10}.  In recent years hydrodynamical simulations have for example been used to quantify the effect of baryons on the DM halo density profiles \\citep[e.g.][]{Gnedin04,Gustafsson06,Duffy10,Tissera10}, spins \\citep[e.g.][]{Bett10,Bryan12}, shapes \\citep[e.g.][]{Abadi10,Kazantzidis04,Read08,Bryan11,Bryan12}, and substructure \\citep[e.g.][]{Dolag09,Romano09} of dark haloes, as well as on the matter power spectrum \\citep[e.g.][]{Jing06,Rudd08,Guillet10,VanDaalen11,Casarini11} and the clustering of subhaloes \\citep{VanDaalen13}. Because the physics of galaxy formation is uncertain, it is important to vary the parameters of the model. In particular, it has recently become clear that the efficient feedback that is required to reproduce observations, and which is thought to be driven by star formation and by AGN at low and high halo masses, respectively, leads to results that differ qualitatively from the predictions of earlier simulations that suffered from overcooling. For example, efficient feedback reduces, or even reverses, adiabatic contraction in the inner parts of massive haloes \\citep[e.g.][]{Duffy10,Mead10,Teyssier11,Killedar12,Martizzi13a}. The ejection of baryons by outflows reduces the matter power spectrum on remarkably large scales \\citep{VanDaalen11} with dramatic consequences for future cosmological weak lensing studies \\citep{Semboloni11,Semboloni12,Zentner13}.  One of the most important quantities characterising the distribution of matter is the halo mass function (HMF hereafter), i.e.\\ the number density of haloes as a function of their mass. The evolution of the HMF, in particular its massive end, is for example a powerful tool for constraining cosmological parameters such as the dark energy equation of state, using future large surveys such as eROSITA \\citep{eRosita12}, XMM-XXL \\citep{Pierre11}, and XCS \\citep{Mehrtens12} in X-ray; Planck\\footnote{http://www.rssd.esa.int/index.php?project=Planck} using the Sunyaev-Zel'dovich effect; and DES\\footnote{http://www.darkenergysurvey.org/}, Euclid \\citep{Euclid11} and LSST\\footnote{http://www.lsst.org/lsst/} using weak lensing. For a given cosmology, the HMF is usually predicted using DM-only simulations \\citep[e.g.][]{Jenkins01,Reed03,Warren06,Lukic07,Tinker08}. To exploit the capacity of upcoming surveys, the theoretical HMF needs to be calibrated at the per cent level \\citep{Wu10}. Therefore, even if the impact of baryon physics is only of the order of a per cent or more, it should be taken into account when computing the theoretical HMF.  The effect of baryon physics on the HMF  has recently been studied using hydrodynamical simulations by, among others, \\citet{Sawala12} and \\citet{Khandai14} at low halo masses and by \\citet{Stanek09}, \\citet{Cui12}, \\citet{Martizzi13b} and \\citet{Cusworth13} at high masses.  Although the studies differ in their detailed findings (and sometimes on the sign of the effect), there is nevertheless a growing consensus that baryon physics will significantly affect the HMF.  Here we will use the suite of cosmological simulations from the OverWhelmingly Large Simulations project (OWLS; \\citealt{Schaye10}) to study the effect of galaxy formation on the mass function and internal structure of haloes more massive than $M^{\\mathrm{crit}}_{200} = 10^{11.5} \\hMsun$. For this purpose we will also make use of new larger volume, lower resolution versions of a subset of the OWLS models (an extension of OWLS, called cosmo-OWLS; see \\citealt{LeBrun14}). OWLS is well-suited to study baryonic effects as it consists of a wide range of models that were run from identical initial conditions, but employing a wide variety of recipes for the uncertain baryonic processes. OWLS also includes a DM-only model as well as a model with AGN feedback that reproduces both optical and X-ray observations of groups and clusters of galaxies \\citep{McCarthy10,LeBrun14}. These last two models are therefore particularly well-suited to our needs and we will employ them to provide fitting formulas that can be used to correct the HMFs predicted by DM-only models for the effect of baryons.  This paper is organised as follows.  In Sec.~\\ref{sec:OWLS} we describe our simulations and explain the methods used to find haloes and to match them between different simulations. In Sec.~\\ref{sec:result} we show how baryon physics, such as radiative cooling and feedback from supernovae and AGN alter the masses of haloes.  In Sec.~\\ref{Sec:fit_formula} we provide analytic fitting formulae to correct the masses of haloes in the DM-only simulation for the effect of baryon physics. In Sec.~\\ref{sec:hmf} we show the impact of baryon physics on the halo mass function and discuss the implications for cluster number count studies. In Sec.~\\ref{sec:discussion} we compare our findings to those of previous studies.  Finally, in Sec.~\\ref{sec:conclusions} we summarize and conclude. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have explored the effects of the introduction of important baryon physics associated with galaxy formation on the total mass of haloes, the mass profile up to large radii ($10R_{200}^{\\rm{dmonly}}$) and the halo mass function (HMF). In order to isolate the effects of baryon physics, we used several simulations from the OWLS project with identical initial conditions, box size and resolution, starting from only gravitationally interacting particles (DMONLY). On top of that we added gas hydrodynamics, star formation and primordial cooling in the NOSN\\_NOZCOOL simulation, an implementation of kinetic SN feedback in the NOZCOOL simulation, introduction of metal-line cooling in the REF simulation and finally the seeding and growth of black holes and AGN feedback in the AGN 8.0 and AGN 8.5 simulations.  We also explored a different prescription for SN feedback in which the mass loading and the wind velocity depend on the local star-forming gas density in the WDENS simulation.  By comparing the results of different simulations, we were able to isolate the importance of different galaxy formation physics on the HMF.   An important aspect of this work was to compare exactly the same set of haloes when different physical processes are introduced in order to isolate the effects of baryons without introducing a bias.  We therefore applied a halo finding and matching algorithm that takes advantage of the fact that the dark matter particles have unique IDs and the simulations all used identical initial conditions.   Using the matched haloes, we compared the relative difference in the total mass of haloes between the DMONLY and the baryonic simulations. Our results span nearly four orders of magnitude in halo mass, $10^{11.5}  < M^{\\mathrm{crit}}_{200} / [\\hMsun]< 10^{15.2} $.  We found that at the low-mass end, SN feedback produces haloes that are 20\\% less massive, at $z=0$, with respect to their DMONLY counterparts, due to the ejection of baryons from the haloes as well as some expansion of the dark matter itself. This difference decreases with increasing halo mass, as the escape velocity gradually becomes too high for the gas to escape, reaching no difference for masses  $M^{\\mathrm{crit}}_{200} > 10^{12.5}\\hMsun$.The mass range over which SN feedback can alter the HMF can be extended somewhat if the wind velocity increases with local gas density (as in WDENS).  However, only AGN feedback can produce a substantial alteration of the HMF of galaxy groups and clusters with halo masses up to $M^{\\mathrm{crit}}_{200} = 10^{14.8}\\hMsun$.   A direct effect of the change in the total mass of the haloes is the modification of the HMF.  Similarly to the total mass variation, we found that supernova feedback is particularly important in shaping the HMF in the mass range $10^{11.5} < M^{\\mathrm{crit}}_{200} < 10^{12.5}\\hMsun $, with a decrease of $20\\%$ with respect to the DM-only simulation. Including only stellar feedback does not produce a significant effect for haloes more massive than $M^{\\mathrm{crit}}_{200} = 10^{13}\\hMsun$.  In the higher mass range, AGN feedback can induce a similar $20\\%$ decrease with respect to the DM-only scenario.  Baryon physics is able to significantly change the total mass profiles of haloes out to several times the virial radius. This means that also the environment in which the haloes reside has significantly different properties with respect to the simulations with only gravitationally interacting particles. This effect could be very important for gravitational lensing measurements, which are sensitive to the mass profile of haloes \\citep[e.g.][]{Mead10,Semboloni11,Killedar12, VanDaalen13}.  We have provided a set of analytic functions that can be used to correct the masses of DMONLY simulated haloes for the presence of baryons (for several mass definitions).  We have shown that the change in mass of the haloes due to baryon physics does not depend on small changes in the values of the cosmological parameters. We also used the analytic fitting formulas to correct the Tinker et al.\\ universal HMF for the effects of baryons.  In this way we are able to predict the abundance of haloes in different cosmologies.  In particular, we showed that the shift in the HMF is about $20\\%$, which has important implication for cluster number counts (e.g., the effect of baryon physics is of the same order as switching the cosmological parameters between WMAP7 and Planck).  To help alleviate this problem, we advocate using only the highest-mass clusters for number counts test, for example $ M^{\\rm mean}_{500} >10^{15}\\hMsun$, where the effect of baryon physics on the mass of the haloes is far less pronounced.  In conclusion we have shown that the masses of haloes inferred from DM-only simulations are not reliable, and when baryon physics is included this can lead to a difference up to $20\\%$ in the mass of the halo and a similar shift in the HMF.  The magnitude of the effect {\\it far} exceeds the percentage precision requirement on the HMF \\citep{Wu10} for future surveys that aim to constrain the dark energy equation of state, such as XMM-XXL, eROSITA, Planck, DES, Euclid and LSST.  Thus, it is beyond question that baryons must be properly modelled for future precision cosmological studies, as well as for any other theoretical studies that require halo masses to be known with better than $\\sim20\\%$ accuracy."
    },
    "1402/1402.0148_arXiv.txt": {
        "abstract": "Building upon a new dynamical model for the X-ray binary LMC X--3, we measure the spin of its black hole primary via the continuum-fitting method. We consider over one thousand thermal-state \\emph{RXTE} X-ray spectra of LMC X--3. Using a large subset of these spectra, we constrain the spin parameter of the black hole to be $\\spin = 0.21^{+0.18}_{-0.22}$ (90\\% confidence). Our estimate of the uncertainty in $\\spin$ takes into account a wide range of systematic errors. The low spin and low mass of the black hole further align LMC X--3 with the class of Roche-lobe overflowing, transient black-hole systems.  We discuss evidence for a correlation between a black hole's spin and the complexity of its X-ray spectrum. ",
        "introduction": "\\label{section:intro}  \\citet{Leong_1971} discovered LMC X--3 during the first year of the {\\it Uhuru} mission. In 1983, \\citet{Cowley_1983} showed via dynamical observations that the compact X-ray source in this 1.7-day binary is a black hole.  In \\citet{Jerry_LMCX3} we use new optical data to derive a much-improved dynamical model of the system, which contains a B-type secondary.  Of chief importance, Orosz et al. report tight constraints on the orbital inclination angle of the binary, $i = 69.6\\pm0.6$~deg, and the mass of the black hole, $M=6.95\\pm0.33~\\Msun$.  LMC X--3 is unusual compared to the full assemblage of black-hole binary systems: On the one hand, like the transient systems, its X-ray intensity is highly variable because the black hole is fed by Roche-lobe overflow.  On the other hand, however, the system almost continually maintains itself in an X-ray bright mode like the persistent (wind-fed) systems \\citep{MNS13, Soria_2001}.   Transient versus persistent black holes are further set apart by the properties of their black hole primaries: the transient black hole masses are low and tightly distributed ($7.8~\\Msun \\pm 1.2~\\Msun$), while the masses of the persistent black holes are appreciably higher ($\\gtrsim11~\\Msun$; \\citealt{Ozel_2010}).  Meanwhile, the spins of the transients are widely distributed ($\\spin \\approx 0-1$), while the three persistent black holes with spin measurements are all high ($\\spin \\gtrsim 0.85$).  The mode of mass transfer firmly identifies LMC X--3 as belonging to the class of transient black hole binaries. The low mass of LMC X--3's primary and its low spin (reported herein) are perfectly congruent with this classification.  Despite strong variations in X-ray brightness ($>3$ orders of magnitude), the X-ray spectrum of LMC X--3 is nearly always in a thermal, disk-dominated state (ideal for measuring spin via the continuum-fitting method), except during occasional prolonged excursions into a low-intensity hard state (e.g., \\citealt{Smale_2012, Wilms_2001}).  Because the X-ray spectrum is strongly disk-dominated, relatively featureless, and minimally affected by interstellar absorption, LMC X--3 is a touchstone for testing spectral models of black-hole accretion disks (e.g., \\citealt{Kubota_2010,Straub_2011}).  In a precursor to this work, we fitted all archival X-ray spectra with reliable flux calibration to a relativistic model of a thin accretion disk.  For hundreds of Rossi X-ray Timing Explorer ({\\it RXTE}) spectra, we showed that the inner radius of the disk is constant to within $\\approx2$\\% over a span of 14 years despite the gross variability of the source.  Furthermore, for an ensemble of eight X-ray missions spanning 26 years, we showed that the radius is constant to better than $\\approx6$\\%, despite the uncertainties associated with cross-calibrating the various detectors. This result is the strongest observational evidence to date that spin can be reliably inferred by measuring the inner-disk radius.  The elegant simplicity of a black hole is encapsulated in the famous \"no-hair theorem'', which tells us that an astrophysical black hole is {\\em completely} described by just its mass and spin angular momentum (the third parameter, electrical charge, being effectively neutral in astrophysical settings).  A spinning black hole is an enormous capacitor of angular momentum, with the spin as a ready energy source that can be tapped mechanically in a black hole's ergosphere.  Spin has long been proposed as the likely energy source behind the enormously energetic jets emitted from black holes (e.g., \\citealt{BZ77}).  This assumption has gained recent support both theoretically (e.g., \\citealt{Tchekhovskoy_2011}), and observationally (\\citealt{NM12,Steiner_jets_2013,MNS13}, but see \\citealt{Russell_2013}).  The two primary means by which black hole spins are measured are 1) X-ray continuum-fitting \\citep{Zhang_1997}, and 2) modeling relativistic reflection (frequently termed the ``Fe-line'' method) \\citep{Fabian_1989}.  The single precept which underpins both methods is the monotonic relationship between spin and the radius of the innermost-stable circular orbit (ISCO) for particles orbiting the black hole.  Consequently, by determining $\\Risco$, which is presumed to be the inner-radius of the accretion disk, one may directly infer a black hole's angular momentum $J$, usually expressed as the dimensionless spin parameter, \\begin{equation} \\spin \\equiv cJ/GM^2, \\qquad  0 \\leq |{\\spin}| \\leq 1. \\label{eq:spin} \\end{equation}  The continuum-fitting method -- the basis for the work presented here -- has been used to estimate the spins of roughly a dozen stellar-mass black holes (e.g., \\citealt{MNS13}, and references therein, \\citealt{Middleton_2006,Mari_GX}).  In this method, the inner-disk radius is estimated using the thermal, multicolor blackbody continuum emission from the disk (e.g., \\citealt{NT73,SS73}).  The Fe-line method measures $\\Rin$ using the redward extent of relativistic broadening of reflection features in the disk, and has been applied just as widely (e.g., \\citealt{Brenneman_Reynolds, Walton_2012a, Reis_2013a} and \\citealt{Reynolds_2013} and references therein). Both methods, but especially the Fe-line method, are also being applied to measure the spins of supermassive black holes in active galactic nuclei (AGN).  LMC X--3 is ideal for continuum-fitting as it offers a strong, dominantly thermal continuum at nearly all times. However, its spectra accordingly contain very little signal in reflection, and so reflection models cannot constrain spin for this source.   Prior to this investigation, a preliminary estimate of LMC X--3's spin was obtained via a continuum-fitting measurement from a single {\\it BeppoSAX} spectrum by \\citet{Davis_2006}, later bolstered by \\citet{Kubota_2010}.  Their results were hampered primarily by the poorly constrained mass available at the time, and so only a rough estimate was possible, $\\spin \\sim 0.3$.  Employing the new and precise mass measurement from \\citet{Jerry_LMCX3}, we revisit LMC X--3's spin determination.  Using the largest data set available to us -- over a decade of pointed monitoring with ({\\it RXTE}) -- we estimate the spin and its uncertainty for LMC X--3 via X-ray continuum fitting. ",
        "conclusions": "\\label{section:concs}  We have analyzed all 1598 spectra of LMC X--3 collected during the {\\it RXTE} mission.  Using a selected sample of $\\approx400$ spectra, our precise measurement of black hole mass and inclination \\citep{Jerry_LMCX3}, and the continuum-fitting method, we derive a strong constraint on the spin of the black hole: $\\spin = 0.21^{+0.18}_{-0.22}$ (90\\% confidence).  Our comprehensive error estimate takes into account all known sources of uncertainty, e.g., the uncertainties in $M$, $i$, $D$, $\\alpha$, and in the absolute X-ray flux calibration.   The simple and predominately thermal spectra of LMC X--3 and A0620--00, the black holes with the lowest measured spins, contrast sharply with the complex and strongly Comptonized spectra of GRS~1915+105 and Cyg X--1, the two black holes with near-extreme spin. This comparison suggests a possible link between spin and the degree of spectral complexity, a hypothesis that can be tested, and which predicts a low spin for GS~2000+25 and a high spin for V404 Cyg.  The black hole 4U 1957+11 may allow a falsification of the hypothesis, if its spin can be measured.  By virtue of the no-hair theorem, we have a complete and quite precise description of the black hole in LMC X--3.  These three -- and only three -- characteristics define the black hole in entirety: a charge of zero, $M \\approx 7.0~\\Msun$ and $\\spin \\approx 0.2$."
    },
    "1402/1402.0181_arXiv.txt": {
        "abstract": "The nonbaryonic dark matter of the Universe is assumed to consist of new stable forms of matter. Their stability reflects symmetry of micro world and particle candidates for cosmological dark matter are the lightest particles that bear new conserved quantum numbers. Dark matter candidates can appear in the new families of quarks and leptons and the existence of new stable charged leptons and quarks is possible, if they are hidden in elusive \"dark atoms\". Such possibility, strongly restricted by the constraints on anomalous isotopes of light elements, is not excluded in scenarios that predict stable double charged particles. The excessive -2 charged particles are bound in these scenarios with primordial helium in O-helium \"atoms\", maintaining specific nuclear-interacting form of the dark matter, which may provide an interesting solution for the puzzles of the direct dark matter searches. ",
        "introduction": "Extensions of the standard model imply new symmetries and new particle states. In particle theory Noether's theorem relates the exact symmetry to conservation of respective charge. If the symmetry is strict, the charge is strictly conserved. The lightest particle, bearing this charge, is stable. It gives rise to the deep relationship between dark matter candidates and particle symmetry beyond the Standard model.  According to the modern cosmology, the dark matter, corresponding to $\\sim 25\\%$ of the total cosmological density, is nonbaryonic and consists of new stable forms of matter. These forms of matter (see e.g. Refs. \\citen{book,newBook,DMRev} for review and reference) should be stable, saturate the measured dark matter density and decouple from plasma and radiation at least before the beginning of matter dominated stage. The easiest way to satisfy these conditions is to involve neutral elementary weakly interacting particles. However it is not the only particle physics solution for the dark matter problem and more evolved models of the physical nature of dark matter are possible.  Formation of the Large Scale Structure of the Universe from small initial density fluctuations is one of the most important reasons for the {\\it nonbaryonic} nature of the dark matter that is decoupled from matter and radiation and provides the effective growth of these fluctuations before recombination. It implies dark matter candidates from the physics beyond the Standard model (see Refs. \\citen{DMRev,LSSFW,Gelmini,Aprile:2009zzd,Feng:2010gw} for recent review). On the other hand, the initial density fluctuations, coming from the very early Universe are also originated from physics beyond the Standard model. In the present review we give some examples, linking the primordial seeds of galaxy formation to effects of particle symmetry breaking at very high energies.  Here we don't touch the exciting problems of the possible nature of dark matter related with extra dimensions and brane cosmology, but even in the case of our 1+3 dimensional space-time we find a lot of examples of nontrivial candidates for cosmological dark matter.  In the Section \\ref{pattern} we present examples of various types of particle candidates for dark matter. We then pay special attention to a possibility for stable charged species of new quarks and leptons to form dark matter, hidden in neutral dark atoms (Section \\ref{asymmetry}). In Section \\ref{ohe} we consider specific form of O-helium (OHe) dark atoms that consist of heavy -2 charged heavy lepton-like particle surrounded by helium nuclear shell. The qualitative advantages of this OHe scenario and the problems of its proof on the basis of a strict quantum mechanical solution of the problem of OHe interaction with nuclei are discussed in Section \\ref{solution}. The conclusive Section \\ref{Discussion} considers the challenges for experimental test of the OHe solution for the puzzles of dark matter searches. ",
        "conclusions": "\\label{Discussion} The existence of heavy stable particles is one of the popular solutions for the dark matter problem. Usually they are considered to be electrically neutral. But potentially dark matter can be formed by stable heavy charged particles bound in neutral atom-like states by Coulomb attraction. Analysis of the cosmological data and atomic composition of the Universe gives the constrains on the particle charge showing that  only $-2$ charged constituents, being trapped by primordial helium in neutral O-helium states, can avoid the problem of overproduction of the anomalous isotopes of chemical elements, which are severely constrained by observations. Cosmological model of O-helium dark matter can even explain puzzles of direct dark matter searches.  The proposed explanation is based on the mechanism of low energy binding of OHe with nuclei. Within the uncertainty of nuclear physics parameters there exists a range at which OHe binding energy with sodium is in the interval 2-4 keV. Annual modulation in radiative capture of OHe to this bound state leads to the corresponding energy release observed as an ionization signal in DAMA/NaI and DAMA/LIBRA experiments.    With the account for high sensitivity of the numerical results to the values of nuclear parameters and for the approximations, made in the calculations, the presented results can be considered only as an illustration of the possibility to explain puzzles of dark matter search in the framework of composite dark matter scenario. An interesting feature of this explanation is a conclusion that the ionization signal may be absent in detectors containing light (e.g. $^3He$) or heavy (e.g. Xe) elements. Therefore test of results of DAMA/NaI and DAMA/LIBRA experiments by other experimental groups can become a very nontrivial task. Recent indications to positive result in the matter of CRESST detector\\cite{cresst}, in which OHe binding is expected together with absence of signal in xenon detectors\\cite{xenon,lux}, may qualitatively favor the presented approach. For the same chemical content an order of magnitude suppression in cryogenic detectors can explain why indications to positive effect in CoGeNT experiment\\cite{cogent} can be compatible with the constraints of CDMS/Ge experiment\\cite{CDMS3}. The model predicts a possibility of OHe binding with silicon, but this effect should be suppressed at low temperature in CDMS/Si experiment\\cite{CDMSi}.  The present explanation contains distinct features, by which it can be distinguished from other recent approaches to this problem\\cite{Edward,Foot,Feng1,Feng2,Drob,Feldstein:2009tr,Bai,Feldstein:2009np,Fitzpatrick:2010em,Andreas:2010dz,Alves:2010dd,Barger:2010yn,Savage:2010tg,Hooper:2010uy,Chang:2010pr,Chang:2010en,Barger:2010gv,Feldstein:2010su}  An inevitable consequence of the proposed explanation is appearance in the matter of underground detectors anomalous superheavy isotopes, having the mass roughly by $m_o$ larger, than ordinary isotopes of the corresponding elements.  It is interesting to note that in the framework of the presented approach positive result of experimental search for WIMPs by effect of their nuclear recoil would be a signature for a multicomponent nature of dark matter. Such OHe+WIMPs multicomponent dark matter scenarios naturally follow from AC model \\cite{FKS} and can be realized in models of Walking technicolor \\cite{KK2}.  Stable $-2$ charge states ($O^{--}$) can be elementary like AC-leptons or technileptons, or look like technibaryons. The latter, composed of techniquarks, reveal their structure at much higher energy scale and should be produced at LHC as elementary species. The signature  for AC leptons and techniparticles is unique and distinctive what  allows to separate them  from other hypothetical exotic particles.  Since simultaneous production of three $U \\bar U$ pairs and their conversion in two doubly charged quark clusters $UUU$ is suppressed, the only possibility to test the models of composite dark matter from 4th generation in the collider experiments is a search for production of stable hadrons containing single $U$ or $\\bar U$ like $Uud$ and $\\bar U u$/$\\bar U d$.  The presented approach sheds new light on the physical nature of dark matter. Specific properties of dark atoms and their constituents are challenging for the experimental search. The development of quantitative description of OHe interaction with matter confronted with the experimental data will provide the complete test of the composite dark matter model. It challenges search for stable double charged particles at accelerators and cosmic rays as direct experimental probe for charged constituents of dark atoms of dark matter."
    },
    "1402/1402.2521_arXiv.txt": {
        "abstract": "\\noindent {\\textit{Context}: Over the past few years, signatures of supernova remnants have been detected in gamma rays, particularly those that are known to be associated with molecular clouds. Whether these gamma rays are generated by cosmic-ray electrons emitting bremsstrahlung or experiencing inverse Compton scattering, or by cosmic-ray protons interacting with ambient hydrogen is usually not known. The detection of hadronic ionization signatures in spatial coincidence with gamma-ray signatures can help to unambiguously identify supernova remnants as sources of cosmic-ray protons. \\\\ \\textit{Aims}: Our central aim is to develop a method to investigate whether the gamma rays are formed by cosmic-ray protons. To achieve this goal, we derived the position-dependent cosmic-ray-induced and photoinduced ionization rates. \\\\ \\textit{Methods}: To calculate hadronic signatures from cosmic-ray-induced ionization to examine the origin of the observed gamma rays from molecular clouds associated with supernova remnants, we solved analytically the transport equation for cosmic-ray protons propagating in a molecular cloud, including the relevant momentum-loss processes, and determined the proton flux at any penetration depth into the cloud. \\\\ \\textit{Results}: Because the solution of the transport equation is obtained for arbitrary source functions, it can be used for a variety of supernova remnants. We derived the corresponding theoretical ionization rate as a function of the penetration depth from the position-dependent proton flux, and compared it with photoinduced ionization profiles for available X-ray data in a case study with four supernova remnants associated with molecular clouds. Three of the remnants show a clear dominance of the hadronically induced ionization rate, while for one remnant, X-ray emission seems to dominate by a factor of 10. \\\\ \\textit{Conclusions}: This is the first derivation of position-dependent profiles for cosmic-ray-induced ionization with an analytic solution for arbitrary cosmic-ray source spectra. The cosmic-ray-induced ionization has to be compared with photoionization for strong X-ray sources. For this purpose, measurements of X-ray spectra from supernova remnant shocks in the sub-keV to keV domain are necessary for a proper comparison. For sources dominated by cosmic-ray-induced ionization (e.g., W49B), the ionization profiles can be used in the future to map the spatial structure of hadronic gamma rays and rotation-vibrational lines induced by cosmic-ray protons. With instruments such as ALMA for the line signatures and CTA for the gamma-ray detection, this correlation study will help to identify sources of hadronic cosmic rays. \\\\ \\\\ \\noindent Keywords: Astroparticle physics -- Radiation mechanisms: non-thermal -- ISM: clouds -- ISM: cosmic rays -- ISM: supernova remnants -- Gamma rays: ISM} ",
        "introduction": "} Supernova remnants (SNRs) are the main candidates for accelerating Galactic cosmic rays (CRs) \\citep[see, e.g.,][]{baade1934,ackermann2013}. Particularly interesting are SNRs associated with molecular clouds (MCs) that are illuminated by CR protons accelerated at the SNR shock front, leading to the emission of gamma rays. Currently, there are about ten of these objects known to emit photons at GeV energies with typical values of $100\\,\\mathrm{MeV}\\lesssim\\,E_\\gamma\\,\\lesssim\\,10\\,\\mathrm{TeV}$ detected with FermiLAT \\citep[see, e.g.,][]{abdo(W51C)2009, abdo(W44)2010} and H.E.S.S.\\ \\citep[see, e.g.,][]{aharonian2008}. Generally, there are three processes that might lead to the production of these gamma rays: Bremsstrahlung or inverse Compton scattering from CR electrons, or the decay of neutral pions that are formed via inelastic scattering of CR protons on ambient hydrogen. If only the detected gamma-ray spectrum is available, the processes causing the gamma radiation cannot be determined definitely from these observations since they neither allow distinguishing a leptonic from a hadronic (proton) scenario, nor do they help telling apart bremsstrahlung from inverse Compton scattering. %  Although sometimes estimates for the magnetic field strength or the seed photon field for inverse Compton scattering, or deductions of average hydrogen densities from observations can be made, there is still plenty of room for interpretation regarding the origin of the gamma radiation. One way to unambiguously recognize pion production as the source of the gamma rays is the detection of neutrinos formed by the decay of charged pions, which are generated via inelastic scattering of CR protons on ambient hydrogen at a fixed ratio to the formation of neutral pions \\citep{particledatabook2010}. A first evidence for astrophysical neutrinos was recently announced by the IceCube collaboration \\citep{icecube2013}. But at this point, the observed flux of extraterrestrial neutrinos is diffuse, and the statistics does not allow statements on possible point sources. Another way to indirectly explore whether the origin of the detected gamma radiation is of hadronic nature is given by the idea that if the gamma rays are induced by high-energy CR protons ($E_{\\mathrm{p}}\\gtrsim1\\,$GeV), one also expects CR protons of lower energy, which, because of the threshold energy for pion production, do not contribute to the gamma flux \\citep{becker2011}. These low-energy protons are quite efficient in ionizing the MCs, even more efficient than CR electrons \\citep{padovani2009}, because of two effects. First, the ionization cross-section for electrons is lower than the one for protons. Second, the lower rest mass of the electrons causes them (a) to be deflected by electromagnetic fields to a higher degree, and (b) to lose their energy more rapidly than CR protons in interactions with matter. Therefore, the CR electrons cannot penetrate an MC deep enough to contribute significantly to the overall ionization rate in its interior. Hence, among all the scenarios proposed to explain the observed gamma radiation, only the hadronic one would imply effective ionization of MCs, particularly in their inner regions. A correlation study of both high-energy gamma rays and low-energy ionization signatures in spatial coincidence might therefore provide a strong support for a hadronic scenario for the generation of gamma rays.  To be able to attribute ionization signatures to low-energy CR protons, other possible ionization sources need to be ruled out, in particular UV photons and X-rays. UV photons are typically absorbed quite effectively already at low column densities of traversed matter, therefore, they play a crucial role in the total ionization only at the outer regions of the clouds \\citep{tielens2010}. X-rays, however, can penetrate the clouds considerably deeper and are also quite effective in their ionization. Consequently, photoionization is the process inside MCs that is required to be exceeded by CR-induced ionization to enable one to attribute detectable ionization signatures to CRs (here, the term \\textit{CR-induced ionization} denotes ionization induced only by primary CR protons and the corresponding secondary electrons). These signatures can be distinguished in terms of their variation in the column-density dependence in the cloud. In this study, the ionization rates induced by both processes are derived as functions of the penetration depth into the cloud, thus providing a crucial clue to the dominant processes that contribute to the formation of gamma rays from SNRs associated with MCs. \\\\ \\\\ The paper is organized as follows: first, the relation between gamma-ray observations and a potential underlying CR proton spectrum is described in Sect.\\ \\ref{CR_prot}. In Sect.\\ \\ref{TE}, the transport equation for the propagation of the CR protons into and inside the MCs is introduced, and both the general analytic solution and a solution for specific choices of the diffusion coefficient, momentum losses, and source function are given. For four specific SNR-MC systems, the position-dependent ionization rates induced by CR protons are derived, using the solution of the transport equation, in Sect.\\ \\ref{CR_ion_prof}. Section \\ref{X-ray} deals with the photon fluxes both at the surfaces and inside the clouds of these systems and the corresponding position-dependent photoionization rates. The ionization rates induced by photons and CR protons are compared in Sect.\\ \\ref{results}. Finally, in Sect.\\ \\ref{conclusions}, we give a brief overview of how ionization rates affect MCs and how they can be detected are provided. We also provide conclusions from the results from Sect.\\ \\ref{results} and give an outlook for future work. ",
        "conclusions": "} We calculated the position-dependent H$_2$ ionization rates of molecular clouds near supernova remnants, induced by photons and cosmic-ray protons, respectively, and compared them to examine a potential hadronic formation scenario of the observed gamma radiation from these systems. The cosmic-ray proton fluxes were constructed by means of the analytic solution of the transport equation for cosmic-ray protons inside the clouds, while the photon fluxes were obtained using the Beer-Lambert law. In addition to the systems consisting of a supernova remnant and a molecular cloud that we studied here, the solution of the transport equation can also be used for many other astrophysical objects with scalar, momentum-dependent diffusion and any type of momentum losses, such as stellar winds, cosmic-ray diffusion in the interstellar medium, or gamma-ray bursts.  Enhanced ionization rates compared with the ionization rates induced by photons alone are expected if the gamma-ray emission seen from systems containing supernova remnants and molecular clouds originates from the decay of neutral pions that are formed in proton-proton interactions of the cosmic-ray protons with ambient hydrogen. The method presented here is particularly promising for objects of low X-ray fluxes and steep, incident cosmic-ray proton spectra, when the spectra can be extrapolated toward lower energies by single power-laws. % For the SNR-MC system W49B, the ionization rate induced by cosmic-ray protons significantly exceeds the ionization rate induced by photons. In contrast, for CTB~37A, the total ionization rate is dominated by photoinduced ionization, while for the objects W44 and 3C~391, the results presented here favor cosmic-ray protons as the dominant source of ionization. The results shown in Figs.\\ \\ref{ion_diff_X} and \\ref{ion_tot_X} are a definite motivation for further studies of enhanced ionization rates attributable to cosmic-ray sources. But since the ionization rate is not directly detectable, one has to derive it for instance from observational spectroscopic data. One method to estimate the ionization rate is to examine the absolute intensity of molecular rotation-vibrational lines, as stated in \\cite{becker2011}. It is also possible to analyze the ratio of the abundance of certain molecules or molecular ions \\citep[see, e.g.,][]{odonnell1974, black1977, indriolo2009}. Computing rotation-vibrational lines requires solid estimates on several physical parameters of the molecular clouds, such as composition, temperature, ionization degree, electron density, and a few others. Many of these parameters are quite difficult to estimate, which means that predictions for rotation-vibrational line intensities are inaccurate. Hence, the molecular ion H$_3^+$ is especially well-suited for observational examinations of ionization rates because only very few cloud parameters are involved in calculating its formation rate \\citep{goto2013}. Particularly in W49B and 3C~391, very high ionization rates are found ($10^{-12}\\,$s$^{-1}$ at depths of 1\\,pc into clouds with a hydrogen density of n$_{\\rm H}=100\\,$cm$^{-3}$). The resulting ionization degrees are on the order of $10^{-4}$ at depths of 1\\,pc and $10^{-3}$ or higher at the cloud surfaces. However, a penetration depth of 1\\,pc at a density of 100\\,cm$^{-3}$ corresponds to a column density of $3\\cdot10^{20}$\\,cm$^{-2}$, so these regions will also be significantly affected by stellar Far UV radiation ($E<13.6\\,$eV), which has an impact on the chemistry and temperature in the clouds. Accordingly, it is also possible that these outer layers of the clouds are photodissociation regions. This needs to be taken into account in the astrochemistry models for the prediction of molecular rotation-vibrational lines to test the theoretical ionization rates by means of observations.  \\begin{table}[htb] \\centering{ \\begin{tabular}{ccccccc} \\hline\\hline instrument        &    H.E.S.S.    & \\multicolumn{2}{c}{FermiLAT}  & CTA & \\multicolumn{2}{c}{ALMA} \\\\ \\hline energy or $\\lambda$ &    >100 GeV    &    0.1 GeV     &    100 GeV     &  10 TeV  &   72.5\\,mm    &    362.5\\,$\\mu$m    \\\\ ang.\\ resolution & 0.1$^{\\circ}$  &  3$^{\\circ}$   & 0.04$^{\\circ}$ &  1' &     1''    &    0.005''    \\\\ object  &        \\multicolumn{6}{c}{linear spatial resolution for 90$^{\\circ}$ inclination [pc]}   \\\\ \\hline W49B    &             14.0      &      419       &      5.6       & 2.3 &   0.0388   &    0.000194   \\\\ W44     &              5.8      &      173       &      2.3       & 1.0 &   0.0160   &    0.000080   \\\\ CTB 37A &             14.0      &      419       &      5.6       & 2.3 &   0.0388   &    0.000194   \\\\ 3C 391  &             12.6      &      377       &      5.0       & 2.1 &   0.0349   &    0.000175   \\\\ \\hline \\end{tabular} \\vspace{0.5cm} \\caption{Spatial resolutions of the studied SNR-MC systems for a selection of relevant telescopes. The distances of the objects from Earth were taken from Table \\ref{tab_dist_ext}. The corresponding relative errors also apply to the spatial resolutions presented here. For ALMA, the maximum base length of 14.5\\,km was used. Sources for the instrumental data: H.E.S.S.: \\cite{HESSresolution}, FermiLAT: {http://www-glast.stanford.edu/instrument.html}, CTA: {http://www.cta-observatory.org/?q=node/26}, ALMA: \\cite{ALMAresolution}.} \\label{resolution}} \\end{table}  \\noindent Some of the molecular species linked to high ionization rates are most effectively and hence solely detectable in absorption. This means that the molecular clouds in which these molecules are to be found have to be somewhat close to Earth, and the sightlines toward these clouds need to feature a bright source in the background for illumination such that the light is absorbed by the molecules in order to trace them. This might restrain the target-selection process of systems composed of a supernova remnant associated with molecular clouds to nearby complexes, that is, a few kpc from Earth, such as W44. These are particularly interesting because the resolution available with current and next-generation telescopes will allow studying their substructure. The currently most promising experiments for the suggested correlation study are the Cherenkov Telescope Array (CTA) and FermiLAT for the detection of the gamma rays, and the Atacama Large Millimeter/submillimeter Array (ALMA) for the detection of the molecular lines. CTA will reach a spatial resolution of $\\sim0.03^\\circ$ \\citep{cta2010,cta2011}. While this is still lower than for submm, IR, or FIR telescopes, CTA measurements can be used to determine regions of enhanced gamma-ray emission that then can be observed with IR telescopes to search for a correlation predicted for hadronic emission scenarios. The spatial resolution for a selection of relevant telescopes for the suggested correlation study is given in Table \\ref{resolution}. Since different observations indicate high ionization rates through the measurement of H$_3^+$ \\citep[see, e.g.,][]{indriolo2010}, it is possible to find the ionization-induced signatures that are spatially correlated with the GeV gamma-ray emission, which will help pinpointing supernova remnants as sources of cosmic-ray protons.  Apart from the specific supernova remnants discussed above, there are currently six additional ones that are associated with molecular clouds. However, for most of these no suitable X-ray data are currently available, because the X-ray fluxes in the corresponding energy ranges are either below the detection limit or are not aligned with the gamma-ray emission regions. For these systems, precise measurements of X-ray spectra by experiments such as SWIFT, XMM-Newton, Suzaku, or Chandra, particularly for energies $E\\sim0.1\\,$keV, can help to estimate the photoionization rate.  To obtain even more accurate theoretical predictions for ionization profiles induced by cosmic rays, additional components of both the cosmic-ray spectrum and the cloud matter can be taken into account. Considering also heavier CR particles in addition to protons for calculating the ionization rate of molecular hydrogen, the model can be applied by adequately adapting the Coulomb-loss coefficient and the ionization cross-section, for example by using the Bethe-Bloch approximation \\citep{bethe1933,padovani2009}. The relative abundances of cosmic-ray helium or heavier nuclei at the acceleration regions are poorly known, however. In addition to the ionization of molecular hydrogen in the clouds, the model can also be used to calculate the ionization rates of other species such as helium by simply replacing the ionization cross-section. This further improves estimates of the ionization degree in the molecular clouds required for the chemistry models. %  To conclude, observable ionization signatures are expected and, combined with gamma-ray observations, correlations that can only be caused by a common primary component \\citep[see, e.g.,][]{juliareview2008} can be revealed."
    },
    "1402/1402.7132_arXiv.txt": {
        "abstract": "Interplanetary dust particles (IDPs) are an important constituent of the earth's stratosphere, interstellar and interplanetary medium, cometary comae and tails, etc. Their physical and optical characteristics are significantly influenced by the morphology of silicate aggregates  which form the core in IDPs. In this paper we reinterpret scattering data from laboratory analogs of cosmic silicate aggregates created by Volten et al. \\cite{volten2007}, to extract their morphological features. By evaluating the structure factor, we find that the aggregates are mass fractals with a mass fractal dimension $d_{m} \\simeq 1.75$. The same fractal dimension also characterizes clusters obtained from {\\it diffusion limited aggregation} (DLA). This suggests that the analogs are formed by an irreversible aggregation of stochastically-transported silicate particles. ",
        "introduction": "\\label{section1}  Fractal geometries provide a description for many forms in Nature such as coastlines, trees, blood vessels, fluid flow in porous media, burning wavefronts, dielectric breakdown, diffusion-limited-aggregation (DLA) clusters, bacterial colonies, colloidal aggregates, etc. \\cite{Mand,Barb,Tam}. They exhibit self-similar and scale-invariant properties at all levels of magnification and are characterized by a non-integer fractal dimension. These features arise because the underlying processes have an element of stochasticity in them. Such processes play an important role in shaping the final morphology, and their origin is distinctive in each physical setting.  Irregular and rough aggregates have also been observed in the astronomical context. Naturally found cosmic dust aggregates, known as interplanetary dust particles (IDPs), are collected in earth's lower stratosphere. They are formed when dust grains collide in a turbulent circumstellar dust cloud such as the solar nebula, and are an important constituent of the interstellar medium, interplanetary medium, cometary comae and tails, etc. Mass spectroscopy analysis of IDPs have revealed that their primary constituents are (i) silicates of Fe, Mg, Al and Ca, (ii) complex organic molecules of C, H, O and N, (iii) small carbonaceous particles of graphite, coal and amorphous carbon and (iv) ices of CO$_{2}$, H$_{2}$O and NH$_{3}$  \\cite{gberg,gberg2,riet2002,tsuch,cuppen}.  Amongst these, there is an exclusive abundance of silicates which aggregate to form particle {\\it cores}. They have been  described as {\\it fluffy, loosely-structured particles with high porosity}.  The  other constituents contribute to the outer covering or the {\\it mantle} and are usually contiguous due to flash heating from solar flares and atmospheric entry \\cite{riet1996}. The core, being deep inside retains its morphology. The latter is believed to have a fractal organization characterized by a fractal dimension, but  this belief is not on firm grounds as yet \\cite{riet1993,riet2004}.  As the core morphology affects the physical and optical characteristics of IDPs, its understanding has been the focus of several recent works \\cite{messenger,min07,henn,rag,botet,nisha}.  Two classes of stochastic fractals are found in nature. The first class is that of {\\it surface fractals} whose mass $M$ scales with the radius of gyration $R$ in a Euclidean fashion, i.e., $M \\sim R^{d}$, where $d$ is the dimensionality. However, the surface area $S$ increases with the radius as $S \\sim R^{d_{s}}$, where $d_{s}$ is the  surface fractal dimension and $d-1\\le d_{s} < d$ \\cite{adj}. Interfaces generated in fluid flows, burning wavefronts, dielectric breakdown and deposition processes are examples of surface fractals. The second class is that of {\\it mass fractals} which obey the scaling relationship, $M \\sim R^{d_{m}}$, where $d_{m}$ is the mass fractal dimension and $1 \\le d_{m} < d$. Examples of mass fractals are DLA clusters, bacterial colonies and colloidal agregates. Further, in many situations, mass fractals are bounded by surface fractals \\cite{Mand,Barb,Tam}.  As a matter of fact, the above mass fractals belong to this class.  There are many unanswered questions in the context of fluffy cores or silicate aggregates of IDPs. For example, are they mass fractals, bounded by surface fractals? What is their mass and surface fractal dimension? What kind of aggregation mechanisms are responsible for this morphology? What are the consequences of fractal organization on the evolution of clusters? In this paper, we provide answers to some of these questions using the {\\it real-space correlation function} $C\\left(r\\right)$ and the {\\it momentum-space structure factor} $S\\left(k\\right)$. Smooth morphologies are characterized by the {\\it Porod law} \\cite{op88,pw09}. The {\\it signature} of fractal domains and interfaces is a {\\it power-law decay} with non-integer exponents in $C\\left(r\\right)$ and $S\\left(k\\right)$. As typical experimental morphologies are smooth on some length scales and fractal on others, the behaviors of  $C\\left(r\\right)$ vs. $r$ and $S\\left(k\\right)$ vs. $k$ are characterized by {\\it cross-overs} from one form to another. We identify these features in laboratory analogs of cores of IDPs  created by Volten et al. using magnesio-silica grains, by reinterpreting their light-scattering data \\cite{volten2007}. We find that these aggregates are mass fractals with a fractal dimension $d_{m} \\simeq 1.75$. The same fractal dimension characterizes diffusion limited aggregation (DLA). We therefore infer that aggregation mechanisms of silicate cores in IDPs are stochastic and irreversible as in DLA.  This paper is organized as follows. In Section~\\ref{section2}, we describe the tools for morphology characterization and their usage to obtain mass and surface fractal dimensions. In Section~\\ref{section3}, we describe the experimental analogs of silicate cores in IDPs and obtain the structure factor from their light scattering data to extract fractal properties. In Section~\\ref{section4}, we present a simulation of the DLA cluster, and the evaluation of its structure factor and the corresponding mass fractal dimension. Finally, we conclude with a summary and discussion of our results in Section~\\ref{section5}. ",
        "conclusions": "\\label{section5}  Interstellar dust particles (IDPs) found in the earth's stratosphere are an important constituent of cosmic matter. These comprise of loosely structured silicate cores or aggregates ensconced in a mantle of organic and carbonaceous compounds \\cite{gberg,gberg2,riet2002,tsuch,cuppen}. The organization and optical characteristics of the IDPs are greatly influenced by the morphology of the core. In this paper, we have re-interpreted scattering data from laboratory analogs of silicate cores in IDPs created by Volten et al. \\cite{volten2007} using the correlation function $C(r)$ and the structure factor $S(k)$. This analysis has provided us a means to quantify characteristics such as the size and texture of these aggregates. The presence of fractal architecture is characterized by power laws with non-integer exponents in the structure factor. We found that the silicate aggregates are mass fractals with a fractal dimension $d_{m} \\simeq 1.75$. This value of $d_m$ is the same as the fractal dimension of aggregates obtained in a diffusion limited aggregation model. We therefore conclude that the aggregates are formed by an irreversible aggregation of stochastically transported silicate particles. We have also studied the effect of density of particles on the fractal dimension. Our observation is that $d_{m}$ approaches the Euclidian dimension ($d = 3$) with increasing density.  Further, we wish to emphasize that $C(r)$ and $S(k)$ contain information averaged over {\\it all} domains and interfaces in contrast to the conventionally used (local) box-counting procedures. Our estimates of $d_{m}$ are therefore very accurate.  For confirmation of our results and further insights, we require data from real cosmic dust and cometary particles. We understand that it is difficult to obtain light scattering data from stellar objects. Alternatively, information providing depth profiles of these assemblies could also be used to evaluate the $C(r)$ and $S(k)$. As discussed above, they are excellent tools for morphology characterization especially due to their direct experimental relevance. More generally, micro-scale phenomena are characterized by mass-dependent diffusion, i.e., the diffusion rate $D(m) \\sim m^{-\\alpha}$, where $m$ is the mass or number of particles in the cluster and $\\alpha$ is a system-specific parameter \\cite{sbp}. The fractal characteristics of aggregates are greatly influenced by $\\alpha$. We are presently investigating them to quantify this influence. We hope that such analyses will enhance our understanding of diffusion mechanisms in mega-scale systems found in the cosmic environment."
    },
    "1402/1402.0524_arXiv.txt": {
        "abstract": "{Cygnus X-1 is the candidate with the highest probability of containing a black hole among the X-ray binary systems in the Galaxy. It is also by far the most often studied of these objects. Recently, the International Gamma-Ray Astrophysics Laboratory Imager onboard then Integral satellite ({\\it INTEGRAL}/IBIS) detected strong polarization in the high-energy radiation of this source, between 400 keV and 2 MeV. This radiation has been attributed to a jet launched by the black hole.} {We consider whether the corona around the black hole might be the site of production of the polarized emission instead of the jet. } {We studied self-consistently the injection of nonthermal particles in the hot, magnetized plasma around the black hole.} {We show that both the high-energy spectrum and polarization of Cygnus X-1 in the low-hard state can originate in the corona, without needing to invoke a jet. We estimate the degree of polarization in the intermediate state, where there is no jet, to provide a tool to test our model.} {Contrary to the commonly accepted view, the jet might not be the source of the MeV polarized tail in the spectrum of Cygnus X-1.} ",
        "introduction": "Cygnus X-1 is a well-studied X-ray binary. The system is composed of a black-hole candidate of  $14.8 M_{\\odot}$ and an early-type O9.7 Iab star of $\\sim 20 M_{\\odot}$ \\citep{orosz2011}. A radio jet was discovered by \\citet{stirling2001} and nonthermal high-energy emission was observed by different instruments, from hundreds of keV to GeV, and even perhaps TeV energies \\citep{McConnell2000,McConnell2002,albert2007,bodaghee2013,sabatini2010,sabatini2013}.  Recently, \\citet{laurent2011} reported the detection and measurement of polarization in the high-energy tail of the emission of Cygnus X-1 detected by the IBIS instrument of the {\\it INTEGRAL} satellite. The linear polarization between $400$ keV and $2$ MeV is $67\\pm30$ \\%. The polarization at lower energies,in contrast, is quite low. This has been interpreted as evidence of a gamma-ray jet in the system \\citep{hardcastle2011}. The MeV flux would be, in this interpretation, synchrotron radiation produced by ultra-relativistic primary electrons accelerated in situ close to the base of the jet \\citep{zdziarski2012}.  In this work we propose a different scenario for the origin of a polarized MeV tail:  secondary leptonic emission induced by nonthermal particle injection in the hot and magnetized corona around the black hole. Proton acceleration is easier to achieve in such an environment. Then, protons interact with the thermal gas and photons injecting pions and pairs by Bethe-Heitler mechanism. The decay of pions injects muons, which in turn results in more pairs. Absorption of gamma-rays from neutral pion decay is another source of electrons and positrons. As a consequence of all these interactions and decays, a population of energetic secondary pairs appears in the corona. These particles cool mainly by synchrotron emission. The emission becomes partially polarized since part of the magnetic field, which is attached to the accretion disk and twisted around the black hole, is ordered.  In this letter we present the results of our calculations of nonthermal particle interactions in the corona of Cygnus X-1, and we provide estimates of the secondary leptonic spectra, the MeV emission, and the polarization for both the low hard (LH) and the intermediate soft (IS) states of the source. ",
        "conclusions": "\\label{discussion}  We estimated the polarization of the emission predicted by a hybrid thermal/nonthermal lepto-hadronic corona model for Cygnus X-1. This provides an alternative explanation of the high polarization observed by \\citet{laurent2011}. The current data correspond only to the LH state, but our predictions will be testable also with data from the IS state. If polarization is found in the IS, the coronal model will be strongly supported. This can be tested in the future using the polarimetric capabilities of the soft gamma-ray detector onboard {\\it ASTRO-H}."
    },
    "1402/1402.2717_arXiv.txt": {
        "abstract": "The exact mechanism behind Type Ia supernovae (SNe-Ia) and the nature of the progenitors is poorly understood, although several theories vie for supremacy.  Their secured importance to the study of cosmology necessitates a resolution to this problem. Observations of nearby SNe-Ia are therefore critical to determine which theory, if any, is the correct one.  SN 2014J discovered at a relatively close 3.5 Mpc in the galaxy M82 provides such an opportunity.  In this paper we give specific predictions for SN 2014J in the context of the Quark Nova Ia (QN-Ia) model.  Predictions include X-ray luminosities just prior and hundreds of days after the explosion, light curve ``glitches\",  neutrino emission, heavy element creation, and  gravitational signatures. ",
        "introduction": "Type Ia Supernovae (SNe-Ia) are thought to be ``standardizable\" candles capable of measuring cosmological distances.  Their utility in this sense has been used over the past two decades to first discover (Riess et al. 1998; Perlmutter et al. 1999) and then study the accelerated expansion of the Universe; the results from which have deeply shaken our perception of nature.  Despite their overwhelming acceptance as ``standardizable\" candles, the mechanism and nature of the progenitors behind the SN-Ia explosion are poorly understood.  We understand from their spectra that SNe-Ia are explosions of carbon/oxygen white dwarfs (COWDs); how they actually explode is still an area of intense research.  Two theories lead the fray, both requiring the accumulation of mass onto the white dwarf (WD) leading to a thermonuclear explosion: the single-degenerate (SD) channel is the accretion onto the WD from a non-degenerate companion (Whelan \\& Iben 1973) and the double-degenerate (DD) channel provides the mass from the merger of two WDs (Iben \\& Tutukov 1994; Webbink 1984).  In the past year a new channel for the explosion of a WD has been presented (Ouyed \\& Staff 2013), involving the explosion of a companion neutron star (NS) as a Quark-Nova Ia (QN-Ia).  Interestingly SNe-Ia formed through this channel are not standardizable using usual methods (Ouyed et al. 2013a) leading to profound implications for cosmology if the majority are formed in this way.  Resolving the progenitor/mechanism debate is therefore of the utmost importance for SNe-Ia to become more accurate tools for measuring cosmological distances, or even for determining if they can be used for such purposes at all.  Each of the aforementioned theories has unique signatures, that if observed would provide strong evidence for their cause.  However, SNe-Ia are difficult to study because they are rare and hence usually discovered at large distances.  The detection of nearby SNe-Ia therefore offers a unique opportunity to study these explosions in exquisite detail and perhaps lead to a resolution of the progenitor/mechanism debate.  The closest modern-era SNe-Ia observed prior to 2014 were SN 1972E in NGC 5253 at $\\sim$ 2.5-8 Mpc  (e.g., Phillips et al. 1992; Sandage \\& Tammann 1975;  della Valle \\& Melnick 1992; Branch et al. 1994; Sandage et al. 1994) and SN 1986G in NGC 5128 at $\\sim$ 3-5 Mpc.  Unfortunately at the time of these SNe, we did not have the space telescope capabilities we have today.  SN 2014J (Goobar et al. 2014) was discovered on Jan 21 2014 by astronomer Stephen J. Fossey in M82 at $\\sim$ 3.5 Mpc, making it one of the closest SNe-Ia observed in the modern era.  With many of the most sophisticated telescopes trained on this object, SN 2014J provides an unprecedented opportunity to shed light on the mechanism and nature of the progenitors behind SN-Ia explosions.  In the coming months and years we expect a deluge of papers on SN 2014J spanning the entire electromagnetic spectrum.  In this paper we present observable predictions for SN 2014J in the context of the QN-Ia model in an attempt to make sense of what may or may not be seen and to provide motivation for upcoming observing proposals. ",
        "conclusions": "The proximity of SN 2014J offers us an unparalleled opportunity to study a SN-Ia  which may reveal clues as to the nature of the progenitors and the explosion mechanism.  The two leading explosion scenarios (SD and DD channels) have recently been joined by a new intriguing possibility; the QN-Ia.  The relative unfamiliarity of the QN-Ia model makes it easy to dismiss.  However, the QN has been successfully applied to a plethora of other astronomical phenomena including SGRs and AXPs (e.g. Ouyed et al 2010), Gamma-ray bursts (e.g. Ouyed et al. 2011) and Superluminous-supernovae (e.g. Ouyed et al. 2012).  In fact the double-humped light curve observed in SN 2009ip and SN 2010mc (modelled and well fit by Ouyed et al. 2013c) was first predicted by the QN model in 2009 (Ouyed et al. 2009), four years prior to its discovery.  Successes put aside, the QN-Ia model does make bold claims that if true could once again alter our perception of nature.  It is therefore imperative that we either rule out this scenario, or  confirm it with observations from near-by SNe-Ia such as SN 2014J.  In this paper we have provided the observer with a list of QN-Ia signatures which, if observed in SN 2014J, would support the existence of  QNe-Ia.  We do note, however, that many of the late-time predictions rely on the existence of a QS.  It is entirely possible that the QS collapses to a BH during the ($^{56}$Ni + $^{56}$Co) decay phase (see point 1 in \\S 4) thereby providing an (albeit unintended) hedge.  If the QS-BH transition occurs at any other time, however, we should see this in the light curve of SN 2014J (e.g. if it occurs immediately, no spin-down energy will be deposited and the light curve will be purely due to ($^{56}$Ni + $^{56}$Co) decay)."
    },
    "1402/1402.4131_arXiv.txt": {
        "abstract": "{Cluster galaxies are the ideal sites to look at when studying the influence of the environment on the various aspects of galaxies' evolution, such as the changes in their stellar content and morphological transformations. In the framework of \\wings, the WIde-field Nearby Galaxy-cluster Survey, we have obtained optical spectra for $\\sim 6000$ galaxies selected in fields centered on 48 local ($0.04< z <0.07$) X-ray selected clusters to tackle these issues. } {By classifying the spectra based on given spectral lines, we investigate the frequency of the various spectral types as a function both of the clusters' properties and of the galaxies' characteristics. In this way, using the same classification criteria adopted for studies at higher redshift, we can consistently compare the properties of the local cluster population to those of their more distant counterparts.} {We describe a method we have developed to automatically measure the equivalent width of spectral lines in a robust way even in spectra with a non optimal signal to noise. Like this, we can derive a spectral classification  reflecting the stellar content, based on the presence and strength of the \\Oii \\ and \\Hd \\ lines.} {After a quality check, we are able to measure 4381 of the $\\sim 6000$ originally observed spectra, in the fields of 48 clusters, 2744 of which are spectroscopically confirmed cluster members. The spectral classification is then analyzed as a function of galaxies' luminosity, stellar mass, morphology, local density and host cluster's global properties, and compared to higher redshift samples (MORPHS and EDisCS). The vast majority of galaxies in the local clusters population are passive objects, being also the most luminous and massive. At a magnitude limit of M$_V<-18$, galaxies in a post--starburst phase represent only $\\sim 11$\\% of the cluster population and this fraction is reduced to $\\sim 5$\\% at  M$_V<-19.5$, which compares to the 18\\% at the same magnitude limit for high-z clusters. ``Normal'' star forming galaxies (\\ec) are proportionally more common in local clusters.} {The relative occurrence of post--starbursts suggests a very similar quenching efficiency in clusters at redshifts in the $0$ to $\\sim 1$ range. Furthermore, more important than the global environment, the local density seems to be the main driver of galaxy evolution in local clusters, at least with respect to their stellar populations content.} ",
        "introduction": "It is now widely accepted that the evolution of galaxies' properties such as morphology, gas and dust content and stellar population properties, has a very strong dependence on the environment where galaxies are found. Galaxies we observe in local clusters and in the field not only present different characteristics, but also the way that led them to be as we observe them now, that is their evolution, is different. By evolution, in this context, we mean the transformations that galaxies undergo both from the morphological aspect and with respect to their stellar population content, as a function of the cosmic time.  A striking example of this environmental dependency is the paucity of blue, late-type galaxies at the centre of local clusters, at odd with what it was found in higher redshift rich clusters where a large population of blue, star-forming galaxies was already discovered more than 3 decades ago \\citep{butcher78a,butcher78b,butcher84}. This is one of the most clear examples of the consequences of the evolutionary effects observable in galaxies in dense environments, and evidences for it have been gathered during the following years \\citep[see, among the others,][]{couch94,lavery97,margoniner00, depropris03,depropris04,haines09}.  Since then, galaxy clusters have proved to be a very important place to study galactic evolution, and were used to probe an extremely wide range of physical conditions, from the dense cores, to the outermost, low-density regions. This was achieved, during the past years, both from a morphological perspective \\citep[see, e.g.,][]{hausman78,dressler80, buote95,dressler97,fasano00,vandokkum01,goto04,sanchezblazques09, poggianti09} and from the point of view of the stellar populations hosted by these galaxies \\citep[e.g.][]{couch87,dressler83,dressler92,dressler99, postman01,dressler13,gladders13}.  Furthermore, evidences are already in place since many years for the fact that the cluster environment acts not only on a ``global'' scale, but also ``locally''. \\cite{prugniel99}, e.g., demonstrated how the presence of young stellar populations is affected by the local environment, showing that galaxies hosting recent star formation tend to prefer the less dense environment of the field or of poor groups. On the other side, the emission--line galaxies fraction was found to decrease towards the centre of dense clusters \\citep[e.g.][]{biviano97,dale01}, proving the importance of the local conditions. Similarly, the projected distance from the cluster's centre was found to be the main factor driving the differences between passive ellipticals and actively star forming spirals \\citep[see][and references therein]{thomas06}, proving the importance of the global environment.  The best approach when trying to understand how blue, actively star forming galaxies so common in high-redshift galaxy clusters, turn into red and passive systems, is probably to focus on those objects which are experiencing a transition in between the two phases. Among these, there are the so called E+A galaxies \\citep{dressler83,zabludoff96}, also named \\ka \\ and \\ak \\ \\citep{dressler99}, or simply ``post--starburst''. Their designation is derived from the characteristics in their spectra, which lack lines in emission and present strong Balmer absorption lines, typical of A-class stars (hence the ``A'' in their names), having a lifetime of some hundreds of Myr, up to about 1 Gyr. For this reason, such a spectral signature is very often used to identify a burst or an episode of enhanced star formation on those time scales.  The presence of strong hydrogen lines in absorption, and the concomitant absence of any emission line in the optical, can be easily achieved as the duration of the two phenomena producing them is different, with emission lines being typically produced from the intense ionizing continuum of very young (i.e. ages less than $\\sim 2\\times 10^7$ years) stars. Such a combination can appear as the result of a star bursting episode being observed few tens of Myr after the star formation has stopped \\citep[e.g.][]{poggianti99}. Alternatively, a truncation in an otherwise continuous, regularly star forming pattern, might produce moderately strong Balmer absorption lines as well \\citep[e.g.][]{newberry90,leonardi96}.  Hence, in most cases, these are galaxies which have recently experienced a quenching of the star formation activity and are on their way to become passively evolving objects. Thus, despite the fact that they are quite rare if compared to the whole galactic population \\citep[see][and references therein]{quintero04}, they may be one of the important keys to address galaxy evolution questions.  E+A galaxies were initially believed to be a prerogative of galaxy clusters, and a number of mechanisms have been proposed to explain their formation, including high speed galaxy--galaxy interactions \\citep[e.g.][]{zabludoff96}, mergers \\citep{bekki01}, gas stripping by ram--pressure from the hot intracluster gas \\citep{tonnesen07,ma08}. The latter in particular, detected in X--rays, is believed to be one of the main actors in a galaxy's life, causing the quenching of the star formation in those systems which are accreting into clusters, through the efficient removal of their gas reservoirs.  Optical spectra contain the information we need in order to study, through the analysis of their integrated light, the stellar content of a galaxy: the stellar mass, the ages of stellar populations, their metallicity and the dust content are all characteristics that can, at least to some degree, be derived from such kind of data.  The \\wings\\footnote[1]{You can refer to the \\wings \\ website for a description of the project and updated details on the survey and its products: \\tt{http://web.oapd.inaf.it/wings/new/index.html}} project \\citep{fasano06}, and its spectroscopic follow--up \\citep{cava09}, is providing the largest set of homogeneous spectroscopic data for galaxies belonging to nearby clusters. Spectral lines are of fundamental importance to characterize the stellar content and properties of a galaxy, providing also a quick and simple way to broadly classify them, according to their star formation history. The reader can refer to \\cite{poggianti99} for an interpretation, from a theoretical point of view, of the meaning of the various spectroscopic classes.  In this paper we will describe the details of a method used to automatically measure the equivalent width (EW) of the most prominent spectral lines in \\wings \\ optical spectra. The method, presented in \\cite{fritz07} for the first time, is able to deal with spectra of low signal-to-noise, and provides a robust estimate of the uncertainty as well. We present and describe the catalogue of rest-frame EW, providing a spectral classification, and we compare the spectral properties of $\\sim 2000$ galaxies to other characteristics such as their luminosity, mass, morphology and to the environment where they reside.  The paper outline is as follows: we briefly present the dataset in Sect.~\\ref{sec:data} and explain and discuss the measurement method in Sect.~\\ref{sec:method}. In Sect.~\\ref{sec:spec} we point out some observational issues and describe the spectroscopic classification criteria, while in Sect.~\\ref{sec:properties} we describe the properties of galaxies in our sample, based on their spectroscopic classification in relation with both other galaxies' properties (such as luminosity, stellar mass, etc.) and with cluster's properties as well. This is followed by a comparison with high redshift cluster galaxies (Sect.~\\ref{sec:highz}), and by a summary of our results (Sect.~\\ref{sec:conclusions}). Details on the EW catalog, together with an example, are given in Appendix \\ref{sec:catalog}.  Throughout this paper we have adopted the usual convention of identifying emission lines with negative values of the EWs, and absorption lines with positive ones, with all the quantities given at restframe. Whenever stellar masses are used, we always refer to a \\cite{salpeter55} IMF, with masses in the $0.15-120$ M$_\\odot$ range, and stellar mass definition n.2, as discussed in \\cite{longhetti09}. Furthermore, we assume a standard $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology, with $H_0=70$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "\\label{sec:conclusions} In this paper we have presented a catalog of equivalent widths of the most prominent spectral lines in the $\\sim 3700$ to $6600$ \\AA \\ range, automatically measured in medium resolution spectra from the \\wings \\ spectroscopic survey. We use these measures to classify galaxies based on the intensity and presence/absence of the \\Hd \\ and \\Oii \\ lines, and we focus our analysis on those objects which are spectroscopically confirmed cluster members. This classification reflects the stellar content of a galaxy, and is widely used as a mean to derive clues on the star formation history both locally and at higher redshifts.  Together with other physical parameters and characteristics of the \\wings \\ galaxies sample derived in previous works such as morphology \\citep{fasano12}, stellar mass and absolute magnitude \\citep{fritz11}, we investigate the properties of the different spectral types, trying to understand how the characteristics of the stellar populations vary as a function of other galactic properties and of both local and global environment, in local clusters.  Having used a classification scheme which is adopted by several other studies at higher redshifts, we can then compare the characteristics of local cluster galaxies to their more distant progenitors.  We can summarize our results as follows: \\begin{enumerate} \\item the local cluster population at magnitudes $M_V \\leq -18$ is dominated by the passive \\k \\ spectral type, followed by \\ec \\ galaxies, even though possibly contaminated  --especially at high luminosity (or stellar mass)-- by emission line from non-stellar origin. About 11\\% of galaxies display a post--starburst spectrum; \\item \\k \\ galaxies are not only the most luminous ones, but they also display the broader range of V-band luminosities. The \\ec's \\ have a similar spread, but a significantly lower ($\\sim 1.5$ magnitudes) peak luminosity; \\item \\k \\ galaxies are also the slowest in clusters (see Fig.\\ref{fig:rad_dist}, right--hand panel). This is consistent with early types being on average more virialized with respect to late types, and with the findings that they were already in place since z$\\simeq 0.6$ \\citep[see, e.g,][]{smail97}; \\item there is a broad correlation between the morphological and the spectroscopical classification: more than 80\\% of the \\k \\ and \\ka \\ spectra are ellipticals or S0s, while the majority of the recently star--forming galaxies [\\ea, \\eb \\ and \\ak \\ ] are spirals. The \\ec \\ class is almost equally populated by early (E+S0) and late (Sp) morphological types. Yet, the fact that most of the spectral types display a full range of morphologies, suggests that the evolution of these two properties is, at least partially, decoupled; \\item no significant correlation is found between the relative number of emission--line galaxies and the cluster general properties, such as their X-ray luminosities or velocity dispersion; \\item the local density, instead, clearly plays a role in the spectral type segregation: a great fraction of passive galaxies is found at the largest values of projected density, while other spectral types are instead more common in less dense environment; \\item there is a fraction of the \\ec \\ population showing substantial differences with respect to the other emission line classes: they display properties more similar to the passive type. We argue that, in this particular case, part of them might be \\k \\ type galaxies where the \\Oii \\ line is not related to star formation, but has instead a different origin; \\item we confirm the lack of local high luminosity post--starbursts, strengthening the idea that these galaxies have undergone a significant evolution. The luminosity function is also significantly different with respect to the one at high redshift; \\item local post--starbursts display the tendency to prefer earlier morphological types, compared to the high redshift ones (i.e. no irregular \\ka \\ and \\ak \\ are found in the \\wings \\ sample); \\item the star formation quenching efficiency is similar between high and low redshift, as long as clusters with velocity dispersions less than $\\sim1000$ km/s are considered, with some differences possibly due to downsizing effects. This strongly suggests RAM-pressure stripping as the main mechanism for halting star formation in infalling galaxies, across the z$\\sim 0-1$ redshift range. \\end{enumerate}  Finally, we briefly describe the EW catalog (see an example of the typical quantities and format in Appendix \\ref{sec:catalog}), which is made publicly available through both the CDS an the Virtual Observatory. More details on this catalog and, in general, on \\wings \\ and its products can be found in Moretti et al. (subm.).  \\noindent \\vspace{0.75cm} \\par\\noindent {\\bf ACKNOWLEDGMENTS} \\par  \\noindent This work made use of Virtual Observatory tools, namely TOPCAT (http://www.star.bris.ac.uk/~mbt/topcat/).\\\\ Facilities: Anglo Australian Telescope (3.9 m- AAT); William Herschel Telescope (4.2 m- WHT).\\\\ 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 \\texttt{http://www.sdss.org/}.\\\\ We are grateful to the anonymous referee, whose comments and remarks helped us to improve the readability of this work."
    },
    "1402/1402.6244_arXiv.txt": {
        "abstract": "Binary neutron star mergers are studied using nonlinear 3+1 numerical relativity simulations and the analytical effective-one-body (EOB) model. The EOB model predicts quasiuniversal relations between the mass-rescaled gravitational wave frequency and the binding energy at the moment of merger, and certain dimensionless binary tidal coupling constants depending on the stars Love numbers, compactnesses and the binary mass ratio. These relations are quasiuniversal in the sense that, for a given value of the tidal coupling constant, they depend significantly neither on the equation of state nor on the mass ratio, though they do depend on stars spins. The spin dependence is approximately linear for small spins aligned with the orbital angular momentum. The quasiuniversality is a property of the conservative dynamics; nontrivial relations emerge as the binary interaction becomes tidally dominated. This analytical prediction is qualitatively consistent with new, multi-orbit numerical relativity results for the relevant case of equal-mass irrotational binaries. Universal relations are thus expected to characterize neutron star mergers dynamics. In the context of gravitational wave astronomy, these universal relations may be used to constrain the neutron star equation of state using waveforms that model the merger accurately. ",
        "introduction": "\\else \\paragraph*{Introduction.--} \\fi \\label{sec:intro} Binary neutron star (BNS) inspirals are among the most promising sources for the advanced configurations of the ground based gravitational wave (GW) detector network~\\cite{Abadie:2010cf}. Advanced configurations of LIGO and Virgo detectors are expected to listen to $\\sim0.4-400$~yr$^{-1}$ events starting from 2016-19~\\cite{Aasi:2013wya}. Direct GW observations will then probe such systems in the near future. In particular, because the late--inspiral--merger phase depends crucially on the stars internal structure, the measurement of the tidal polarizablity parameters from GWs will put the strongest constraints on the unknown nuclear equation of state~(EOS)~\\cite{Read:2009yp,Damour:2012yf,Read:2013zra,DelPozzo:2013ala}.  An accurate modeling of neutron star mergers requires numerical relativity (NR). In recent years simulations have become fairly robust, but exploring the physical parameter space remains a challenge out of reach. Furthermore, the interpretation of simulation data can be nontrivial: meaningful quantities must be gauge invariant and possibly have well-defined post-Newtonian (PN) limits. The GW phasing analysis for multi-orbits ($\\sim 10$) simulations was performed by some groups, e.g.~\\cite{Baiotti:2010xh,Bernuzzi:2011aq,Hotokezaka:2013mm}. The BNS dynamics, expressed via the gauge-invariant relation between binding energy and angular momentum~\\cite{Damour:2011fu,Taracchini:2013rva}, was recently analyzed in both the nonspinning and spinning case~\\cite{Bernuzzi:2012ci,Bernuzzi:2013rza}. For both observables a solid analytical framework, although approximate, is essential for extracting information from the simulations.  Despite these detailed studies, simple, and fundamental questions about the merger physics still lack of quantitative answers. For instance, a test-mass in the Schwarzschild metric of mass $M$ has a last stable orbit (LSO) at $R_{\\rm LSO}=6M$, (we use units with $G=c=1$) with dimensionless (or mass-reduced) orbital frequency $M\\Omega_{\\rm LSO}^{\\rm Schw}=6^{-3/2}\\approx 0.06804$. The associated GW frequency $2 M\\Omega_{\\rm LSO}^{\\rm Schw}\\approx 0.13608$ is commonly used to mark the end of the quasiadiabatic BNS inspiral, setting $M$ equal to the total mass of the binary. Similarly,  the specific LSO binding energy $E_{b\\, \\rm LSO}^{\\rm Schw}=(8/9)^{1/2}-1\\approx-0.0572$ is used to estimate the total amount of GW energy emitted during the coalescence process. These numbers appear ubiquitously in BNS-related studies, e.g.,~\\cite{DelPozzo:2013ala}, but are, in principle, no more then an order of magnitude estimates as they neglect both finite mass ratio and finite size effects. Some questions arise: How to model/include these effects? How does the merger frequency and binding energy depend on the main parameters of the binary (EOS, mass ratio and individual spins)? How accurate are the Schwarzschild LSO estimates? In this work we use new multi-orbit NR data and the analytical effective-one-body (EOB) approach problem to put forward some answers. We find that the GW frequency and binding energy at the moment of merger are characterized only by certain dimensionless tidal coupling constants (a fact also empirically observed in~\\cite{Read:2013zra} for the frequency) and the stars spins as a consequence of a fundamental property of the underlying conservative dynamics.    \\ifusesec ",
        "conclusions": ""
    },
    "1402/1402.1378_arXiv.txt": {
        "abstract": "As part of a larger search of \\textit{Wide-field Infrared Survey Explorer} (WISE) data for cool brown dwarfs with effective temperatures less than 1000 K, we present the discovery of three new cool brown dwarfs with spectral types later than T7.  Using low-resolution, near-infrared spectra obtained with the NASA Infrared Telescope Facility and the \\textit{Hubble Space Telescope} we derive spectral types of T9.5 for \\WISEzeroninefourthreeFull, T8 for \\WISEtwozerozerozeroFull, and Y0: for \\WISEtwentytwozeronineFull.  The identification of \\WISEtwentytwozeronineFull\\ as a Y dwarf brings the total number of spectroscopically confirmed Y dwarfs to seventeen.  In addition, we present an improved spectrum (i.e. higher signal-to-noise ratio) of the Y0 dwarf \\WISEzerofourtenFull\\ that confirms the Cushing et al.\\ classification of Y0.  Spectrophotometric distance estimates place all three new brown dwarfs at distances less than 12~pc, with \\WISEtwozerozerozeroFull\\ lying at a distance of only 3.9--8.0~pc. Finally, we note that brown dwarfs like \\WISEtwozerozerozeroFull\\ that lie in or near the Galactic plane offer an exciting opportunity to measure their mass via astrometric microlensing. ",
        "introduction": "The launch of the \\textit{Wide-field Infrared Survey Explorer} \\citep[\\WISE;][]{2010AJ....140.1868W} in late 2009 ushered in a new era in the study of brown dwarfs with spectral types later than T8 (effective temperatures (\\teff) less than roughly 700~K).  Previous searches in the field with the Two Micron All Sky Survey \\citep[2MASS;][]{2006AJ....131.1163S,2002ApJ...564..421B,2004AJ....127.2856B}, the UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{2007MNRAS.379.1599L,2008MNRAS.390..304P,2010MNRAS.406.1885B} and the Canada France Brown Dwarf Survey \\citep[CFHBD;][]{2008A&A...484..469D,2011AJ....141..203A} uncovered ten brown dwarfs with spectral types later than T8 but the survey sensitivity limits combined with their intrinsic faintness at red optical and near-infrared wavelengths \\citep[$M_H$ $\\gtrsim$ 19~mag;][]{2013ApJ...762..119M} limited the ability of these surveys to identify large numbers of them.  In contrast, \\WISE\\ surveyed the entire sky at four mid-infrared wavelengths centered at 3.4 $\\mu$m, 4.6 $\\mu$m, 12 $\\mu$m, and 22 $\\mu$m (hereafter denoted $W1$, $W2$, $W3$, and $W4$ respectively) that are ideally suited to detect cool brown dwarfs because these bandpasses span the peak of their spectral energy distributions.  In particular, the $W1$ and $W2$ bands are centered on a deep CH$_4$ absorption band at 3.3 $\\mu$m and a region relatively free of opacity at 4.7 $\\mu$m where a large fraction of the total flux of cool brown dwarfs emerges \\citep[see Figure 2 of][]{2011ApJ...726...30M}.  The resulting $W1$$-$$W2$ color is very red which makes the identification of cool brown dwarf candidates relatively straightforward.  To date, \\WISE\\ data have been used to identify roughly two hundred brown dwarfs with spectral types later than T6 \\citep{2011ApJ...726...30M,2011ApJ...735..116B,2011A&A...532L...5S,2011ApJS..197...19K,2011ApJ...743...50C,2012ApJ...753..156K,2012ApJ...759...60T,2012ApJ...760..152L,2013ApJS..205....6M,2013ApJ...767L...1L,2013PASP..125..809T}, including the sixteen spectroscopically confirmed Y dwarfs previously published \\citep{2011ApJ...743...50C,2012ApJ...753..156K,2012ApJ...759...60T,2012ApJ...758...57L,2013ApJ...776..128K}.   The detailed study of these cool brown dwarfs promises to shed light on the physics of ultracool atmospheres and constrain brown dwarf formation theories.  In particular, simulations by \\citet{2004ApJS..155..191B} suggest that the space density of cool brown dwarfs is very sensitive to the underlying mass function, a potential constraint on the yet unknown formation mechanism(s) of brown dwarfs \\citep[see][for a review of the current theories]{2013arXiv1302.3954S}.  However, brown dwarfs obey a mass-luminosity-age relation so measuring the brown dwarf present-day mass function is difficult given the ages of the vast majority of brown dwarfs are unknown.  The underlying mass function can be constrained, however, by first forward modeling the evolution of a population of brown dwarfs over time \\citep[e.g.,][]{2004ApJS..155..191B,2005ApJ...625..385A,2006MNRAS.371.1722D} and then comparing the observed space densities of brown dwarfs to the predictions of the models.  Given the relative ease with which effective temperatures can be estimated for cool brown dwarfs, the most common technique has been to compare the space densities of the brown dwarfs as a function of spectral type to model effective temperature distributions \\citep[e.g.,][]{2008MNRAS.390..304P,2010MNRAS.406.1885B,2012ApJ...753..156K,2013MNRAS.430.1171D,2013MNRAS.433..457B}. Results to date suggest that if the mass function is modeled as dN/d$M$ $\\propto$ $M^{-\\alpha}$, then $-1$ $<$ $\\alpha$ $<$ 0.   Implicit in this technique is the necessity of a complete, volume-limited sample of brown dwarfs.  In this paper, we report three new discoveries from our on-going program to construct just such a volume-limited sample of cool brown dwarfs using data from \\textit{WISE}.  In addition, we present an improved $J$- and $H$-band spectrum of the the Y0 dwarf \\WISEzerofourtenFull\\ (hereafter \\WISEzerofourten) located at a distance of roughly 4 pc \\citep{2013ApJ...762..119M}. ",
        "conclusions": "The classification of \\WISEtwentytwozeronine\\ as a Y dwarf brings the total number of \\textit{spectroscopically} confirmed Y dwarfs to seventeen\\footnote{WD 0806$-$661 \\citep{2011ApJ...730L...9L} and CFBDSIR J1458$+$1013B \\citep{2011ApJ...740..108L} have estimated effective temperatures similar to those of the known population of Y dwarfs but spectra of these cold brown dwarfs have yet to be obtained due to their faintness.}.  The colors of the late-type T dwarfs and Y dwarfs exhibit distinct changes at the T/Y boundary with the $J-H$ colors plateauing (albeit with a large scatter) \\citep{2011ApJ...743...50C,2012ApJ...753..156K} and the $z'-J$, $z'-H$, and $Y-J$ colors rapidly becoming bluer \\citep{2012ApJ...758...57L,2013A&A...550L...2L}.  Figure \\ref{fig:nirvssptypes} shows the MKO $J-H$ colors of a sample of T and Y dwarfs \\citep{2010ApJ...710.1627L,2013ApJ...763..130L} as a function of spectral type.  The $J-H$ colors plateau in the late-T spectral types and turn redward at later types (with the exception of the Y0.5 dwarf \\WISEfifteenfortyoneFull).  The three new brown dwarfs are shown as filled red circles and fall in line with the known sequence suggesting that they are not peculiar.  The identification of three new late-type brown dwarfs that lie within roughly 10 pc of the Sun is an important addition to the sample of nearby cold brown dwarfs.  We do, however, defer updating the space density estimates of the late-type T dwarfs and Y dwarfs presented in until the parallaxes of these brown dwarfs have been measured and more nearby cool brown dwarfs are identified.   One of the three new brown dwarfs, \\WISEtwozerozerozero, is located in a region ($l$=72$\\fdg$7, $b$=$+$3$\\fdg$32) covered by the UKIDSS Galactic Plane Survey \\citep{2008MNRAS.391..136L}.  \\citet{2010MNRAS.408L..56L} searched the Sixth Data Release of the Galactic Plane Survey to search for late-type T or Y dwarf candidates but identified only a single candidate - the well-known T9 dwarf \\UGPSzeroseventwotwoFull.  The position of \\WISEtwozerozerozero\\ was not covered by the Sixth Data Release but even if it had been, \\WISEtwozerozerozero\\ would have been missed by \\citet{2010MNRAS.408L..56L} because its values of \\texttt{k\\_1Ell} and \\texttt{jppErrbits}, two flags in the UKIDSS catalog that describe the source ellipticity and photometric quality, fall outside of the \\citeauthor{2010MNRAS.408L..56L} search criteria. This suggests that perhaps more brown dwarfs in the Galactic plane await discovery.   Indeed \\WISEtwozerozerozero\\ joins a growing list of brown dwarfs identified in or near the Galactic plane including \\UGPSzeroseventwotwoFull\\ and Luhman's binary system WISE J104915.57$-$531906.1 \\citep{2013ApJ...767L...1L}.  Previous searches for brown dwarfs including our own ignore either part of, or the entire, Galactic plane in order to minimize the impact of stellar contamination \\citep[e.g.,][]{2002ApJ...564..421B,2007AJ....133..439C}.  However, once a brown dwarf has been discovered in the plane, the reasonably high surface density of background stars offers us an opportunity to measure the mass of the brown dwarf via astrometric microlensing \\citep{1964MNRAS.128..295R,1995AcA....45..345P,1996ApJ...470L.113M}.  \\begin{deluxetable}{lccc} \\tablecolumns{11} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablecaption{\\label{tab:properties}Derived Properties of New Brown Dwarfs} \\tablehead{ \\colhead{Object} & \\colhead{Spectral} & \\colhead{$M_H$} & \\colhead{$W2$} \\\\ \\colhead{} & \\colhead{Type} & \\colhead{Distance} & \\colhead{Distance}}   \\startdata  \\WISEzeroninefourthree\\ & T9.5 & 7.9--34.3 & 6.6--12.1 \\\\ \\WISEtwozerozerozero\\ & T8 & 3.9--7.3 & 6.4--8.0 \\\\ \\WISEtwentytwozeronine\\ & Y0: & $\\cdots$ & 5.2--11.7 \\\\  \\enddata  \\end{deluxetable}     During a microlensing event, the lens (i.e. the brown dwarf) both magnifies the source (i.e. the background star) and shifts its centroid. If the angular separation between the lens and the source is much larger than the Einstein radius, then the magnification is negligible but the centroid shift remains \\citep[e.g.,][]{1996ApJ...470L.113M,2001PASP..113..903G}.  If the lens is also dark (i.e. is much fainter than the source) then a measurement of the centroid shift of the source is relatively straightforward and yields the mass of the lens.  Cool T and Y dwarfs are ideal targets for astrometric microlensing measurements at optical wavelengths because they are effectively dark at these wavelengths \\citep[$M_V$ $\\gtrsim$ 26 mag,][]{2010A&A...510A..99K,2013arXiv1312.1303B}.  The expected number of lensing events per year, which is derived in the Appendix, is given by:  \\begin{equation} N = 2\\kappa \\pi_L \\mu_L \\sigma_\\mathrm{bg} \\left ( \\frac{\\delta M_L}{\\delta \\Delta \\theta} \\right ), \\end{equation}  \\noindent where $\\kappa \\equiv \\frac{4G}{c^2\\, \\mathrm{AU}} \\approx 8.144\\, \\frac{\\mathrm{mas}}{M_\\odot}$, $\\pi_L$ is the parallax of the lens in units of arcseconds, $\\mu_L$ is the proper motion of the lens in units of arcseconds per year, $\\sigma_\\mathrm{bg}$ is the surface density of stars in units of stars per square arcseconds, $\\delta M_L$ is the required precision of the mass determination in solar masses, and $\\delta \\Delta \\theta$ is the precision of the astrometric measurement in milliarcseconds.  Although we cannot compute a rate for \\WISEtwozerozerozero\\ because its parallax and proper motion are unknown, we can get a sense of the expected rate of brown dwarf lensing events near the Galactic plane by considering a hypothetical brown dwarf located at $l=45^\\circ$, $b=0^\\circ$, at a distance of 10 pc, and with a proper motion of $0\\farcs$5 yr$^{-1}$.  We used the Basan\\c{c}on model of the Galaxy \\citep{2003A&A...409..523R} to generate a simulated population of stars in a 0$\\fdg$1 $\\times$ 0$\\fdg$1 field at this position which results in a surface density of stars with $V<$ 24 of 0.12 stars per square arcsecond.  If we require a mass precision of $\\delta M_L$=0.10 $M_\\odot$ and assume we can achieve an astrometric precision of $\\delta \\Delta \\theta$=0.2 mas (e.g., with the \\textit{Hubble Space Telescope}), then we find that $N\\approx$ 0.05 or roughly one event every two decades.  The potential to measure the masses of even a handful of cool field brown dwarfs on a decades timescale should therefore stoke interest in searches for brown dwarfs in the zone of avoidance.  \\begin{figure} \\centerline{\\hbox{\\includegraphics[width=3.5in,angle=0]{fig5.eps}}} \\caption{\\label{fig:nirvssptypes}MKO-NIR $J-H$ colors of T and Y dwarfs as a function of spectral type.  The filled black circles are from \\citet{2010ApJ...710.1627L,2013ApJ...763..130L} while the red filled circles are from this work.  \\WISEeighteentwentyeightFull\\ is plotted separately because its spectral type is currently $\\geq$Y2.  The three new brown dwarfs fall in line with the sequence suggesting they are not peculiar.} \\end{figure}     Computing values for \\UGPSzeroseventwotwoFull\\ and Luhman's binary in a similar fashion we find 0.02 and 0.3, respectively.  The value for Luhman's binary is remarkably high, one event every few years, because of its close distance ($\\sim$ 2 pc) and high proper motion (2$\\farcs$7 yr$^{-1}$).  Followup observations to determine whether it will actually lens a nearby star should therefore be a high priority since a mass measurement of such a nearby (binary) brown dwarf would be invaluable."
    },
    "1402/1402.4241_arXiv.txt": {
        "abstract": "{Using over 89 million astrometric observations of 349\\,737 numbered minor planets, an attempt was made to determine the masses of 230 of them by simultaneously solving for corrections to all orbital elements and the masses. For 132 minor planets an acceptable result was obtained, 50 of which appear to be new.} ",
        "introduction": "The astrometric determination of asteroid masses is the technique of calculating the mass of asteroids (hereafter called ''perturbing asteroids'' or ''perturbers'') from the gravitational perturbations they exert on the motion of other asteroids (called ''test asteroids'').  The usual procedure to calculate the mass of one asteroid is, first, to identify potential test asteroids by searching for close approaches with the perturbing asteroids and then computing some parameter (usually the deflection angle or the change in mean motion) characterising the effect of the perturbing asteroid, see e.g. Gal\\'ad and Gray~(\\cite{Gala2002}). Promising cases are then further investigated to see if they yield meaningful results, since the outcome also depends on the number, distribution and quality of the available astrometric observations of the test asteroid. Asteroids not making close approaches but moving in near-commensurable orbits can also be useful for mass determinations (Kochetova~\\cite{Koch2004}).  A better but more laborious way to find suitable close encounters is to integrate the orbits of test asteroids over the period of available observations with and without the perturbations of the perturbing asteroid (Michalak~\\cite{Mich2000}). Significant differences in the geocentric coordinates between the two calculations indicate cases that should be investigated further.  Initially, asteroid masses were computed from their perturbations on one other asteroid, the first instance being the mass of 4~Vesta computed from the orbit of 197~Arete by Hertz~(\\cite{Hert1968}). When several mass values for the same asteroid are calculated from a number of test asteroids, the usual practice is to compute a weighted mean along with the associated uncertainty as the final value, as in Michalak~(\\cite{Mich2000}). Mass determinations where one or a few masses are determined simultaneously from a number of test asteroids have been attempted on a modest scale in the past--see Sitarski and Todororovic-Juchniewicz~(\\cite{Sita1995}), Goffin~(\\cite{Goff2001}) and others.  This paper gives the results of an attempt to simultaneously determine a large number of asteroid masses (230) from a substantial number of test asteroids (349\\,737). Sect. 2 explains the basic principles and problems of simultaneous asteroid mass determination and in Sect. 3 the details of the actual calculations are given. Sect. 4 briefly deals with the observations used and explains the statistical analysis of the residuals that was used to assign weights to, or reject observations. In Sect. 5 the results are presented and Sect.6 discusses them.  Throughout this paper, asteroid masses are expressed in units of $10^{-10}$ solar masses ($M_{\\sun}$) and generally given to 3 decimal places. ",
        "conclusions": "The masses given in Table~\\ref{Masses} are presented ''as such'', i.e. as the outcome of the least-squares adjustment described above. The bulk density is a quantity derived from the mass and the assumed diameter and its uncertainty is thus affected by the uncertainties of both the mass and the diameter. No physical interpretation of the densities will be attempted here. Some results conform to what is to be expected for their taxonomic class, while others are higher or lower. One should bear in mind that the taxonomic class is derived from the spectral properties of the surface of an asteroid and as such does not say much about its internal composition or porosity.  The mass and size of 4~Vesta is accurately known from the orbit of the Dawn spacecraft. Russell et al.~(\\cite{Russ2012}) give a value of $2.59076 \\pm 0.00001 \\times 10^{20}$~kg corresponding to $1.30260 \\pm 0.00001 \\times 10^{-10} M_{\\sun}$, an equivalent diameter of 523.15~km and a density of $3.456~\\mathrm{g/cm^3}$. The result from this study, to 4 decimal places, is $1.2995 \\pm 0.0015$ yielding a density of $3.448~\\mathrm{g/cm^3}$, in good agreement with the Dawn result.  It would be a constructive achievement to develop an objective criterion to exclude test asteroids that do not significantly contribute to the solution. The correlations between the masses and the semi-major axes come to mind as an obvious choice. Using the sum of the absolute values of the $M$~correlation coefficients as an indicator for the usefulness of a test asteroid, test calculations were made with $N = 10,000$ and $M = 10$ (the 10 largest masses). With a limit for the sum of 0.01, about 15\\% were excluded and the average difference in the 10~masses was 0.0003; this became 48\\% and 0.0012 with 0.02 as the limit. It seems worthwhile to test this criterion on the full set of test asteroids.  In general, the two test calculations clearly show that the most influential factor is the treatment of the observations (weighting and outlier rejection). In that respect, it should be remarked that the observations reduced with the USNO~A1.0, A~2.0 and B~1.0 star catalogues have not been corrected for the systematic biases as described in Chesley et al.~(\\cite{Ches10}) or in R\\\"oser et al.~(\\cite{Rose10}). It should be worthwile to include these in future work.  The planetary ephemeris uncertainties are negligible in this case, certainly the positional errors. The planetary masses come from spacecraft encounters, and DE421/422 is up-to-date. The best asteroid-produced planetary mass of Jupiter, for example, was inaccurate by about 6 parts in $10^6$; the spacecraft-produced values are two orders of magnitude more accurate.  The Gaia satellite, launched at the end of 2013, will in the course of its 5 year nominal duration systematically observe $\\sim$300\\,000 asteroids brighter than $V = 20$ with an accuracy of 0.3--5 mas. This will enable the determination of the masses of about 150 individual asteroids from close encounters, $\\sim$100 with an accuracy of $\\le 50\\%$ and $\\sim$50 with an accuracy of $\\le 10\\%$ (Mouret et al.~\\cite{Mour2008}). The Gaia star catalogue will enable the reduction of future, and the re-reduction of old observations with a very high level of accuracy (about 10~mas, see Desmars et al.~\\cite{Desm2013}), thus considerably increasing the accuracy of simultaneous asteroid mass determinations.  Finally, Table~\\ref{Overview} gives an overview of mass results from recent determinations for those minor planets given in Table~\\ref{Masses}, and published in roughly the last ten years. For a more complete list inlcuding older determinations, the reader is referred to Zielenbach~(\\cite{Ziel2011}). Besides dynamical determinations using asteroids as test objects, there are also determinations from natural satellites and from planetary ephemerides. Note that these use a completely different type of observations."
    },
    "1402/1402.6558_arXiv.txt": {
        "abstract": "{} {We investigate the properties of orphan penumbrae, which are photospheric filamentary structures observed in active regions near polarity inversion lines that resemble the penumbra of regular sunspots but are not connected to any umbra.} {We use Hinode data from the Solar Optical Telescope to determine the properties of orphan penumbrae. Spectropolarimetric data are employed to obtain the vector magnetic field and line-of-sight velocities in the photosphere. Magnetograms are used to study the overall evolution of these structures, and G-band and \\ion{Ca}{II}~H filtergrams are to investigate their brightness and apparent horizontal motions. } {Orphan penumbrae form between regions of opposite polarity in places with horizontal magnetic fields. Their magnetic configuration is that of $\\Omega$-shaped flux ropes. In the two cases studied here, the opposite-polarity regions approach each other with time and the whole structure submerges as the penumbral filaments disappear. Orphan penumbrae are very similar to regular penumbrae, including the existence of strong gas flows. Therefore, they could have a similar origin. The main difference between them is the absence of a ``background'' magnetic field in orphan penumbrae. This could explain most of the observed differences.} {The fast flows we detect in orphan penumbrae may be caused by the siphon flow mechanism. Based on the similarities between orphan and regular penumbrae, we propose that the Evershed flow is also a manifestation of siphon flows.} ",
        "introduction": "\\label{introduction}  All structures seen in the solar photosphere other than granulation are caused by magnetic fields that influence the convection and thus the presence and shape of granules. The largest, strongest, and most complex magnetic fields are located in active regions. The fields can completely inhibit convective motions there, and thus create sunspots and pores \\citep[see][]{Keppens:1996, Solanki:2003, Schlichenmaier:2009, Borrero:2011}. There are other structures observed in active regions that are associated with the submergence or emergence of magnetic flux tubes but not so thoroughly studied as the more prominent sunspots and pores.  In this paper, we analyse filamentary structures observed in active regions near polarity inversion lines. Since they visually resemble a penumbra without an umbra, we call them ``orphan penumbrae'' following \\citet{Zirin:1991}. Similar structures were described before.  There are numerous reports of alignments of granules or intergranular lanes near polarity inversion lines of active regions \\citep[first by][]{Miller:1960}. \\citet{Loughhead:1961} interpreted the alignments as being due to magnetic loops rising through the solar photosphere. The same conclusion was reached by other authors, although the reported lifetimes and sizes of the structures differ considerably \\citep{Brants:1985, Tarbell:1990, Wang:1992, Schlichenmaier:2010}. In a recent work, \\citet{Kuckein:2012} suggested that orphan penumbrae are the photospheric signatures of low-lying flux ropes. However, as pointed out by \\citet{Lim:2013} and \\citet{Zuccarello:2013}, other mechanisms can also lead to the formation of orphan penumbrae.  These authors observed orphan penumbrae in regions with overlaying chromospheric fields and interpreted them as a consequence of rising $\\Omega$-loops that are trapped in the photosphere by the latter.  Here, we suggest that the $\\Omega$-loops responsible for orphan penumbrae submerge again during their decay phase.  The presence of inclined magnetic fields in the photosphere may indeed explain the formation of orphan penumbrae from at least a theoretical perspective. \\citet{Chandrasekhar:1961} showed that convection in inclined magnetic fields is modulated by the strength of the vertical field component, while the horizontal component determines the shape of the convective cells. This was observationally supported by \\citet{Jurcak:2011}, assuming that penumbral filaments are of convective origin. Numerical simulations also confirm the importance of the horizontal component of the magnetic field for the convective cell shape \\citep[see][]{Thomas:2008}. If the magnetic field is horizontal, convection will not be stopped but the cells will be highly elongated and resemble filaments, which perhaps, explains the existence of orphan penumbrae.  Here, we determine the physical properties of orphan penumbrae during their decay phase using spectropolarimetric measurements, magnetograms, and G-band and \\ion{Ca}{II}~H filtergrams taken by the Hinode satellite.  Section~\\ref{observations} describes the observations and data analysis.  The magnetic and dynamic configuration of orphan penumbrae is inferred and compared with that of regular penumbrae in Sect.~\\ref{results}.  The evolution, brightness, and horizontal motions of these structures are also discussed there. We interpret the various observational results in Sect.~\\ref{discussion} and summarise our conclusions in Sect.~\\ref{conclusions}.  \\begin{figure*}[!t] \\centering \\includegraphics[width=0.96\\linewidth]{figure1_time_axis.eps} \\caption{Chart showing the observations of active region NOAA~10960. The blue and green dashes indicate BFI observations in G-band and \\ion{Ca}{II}~H, respectively. The red dashes indicate spectropolarimetric scans. The arrows mark the data shown and discussed in this paper.}  \\label{time_axis_2007} \\end{figure*}  \\begin{figure*}[!t] \\centering \\includegraphics[width=0.96\\linewidth]{figure2_time_axis.eps} \\caption{Analogous to Fig.~\\ref{time_axis_2007} but for active region NOAA~11089. The yellow colour indicates NFI V/I~magnetograms in the \\ion{Na}{I}~D. All the NFI V/I~magnetograms are used in this paper, so they are not marked individually by arrows.} \\label{time_axis_2010} \\end{figure*} ",
        "conclusions": "\\label{conclusions}  We thoroughly investigated orphan penumbrae, which are filamentary structures resembling sunspot penumbrae but are not connected to an umbra that lie near PILs in active regions.  We observed these structures in two active regions and used Hinode SOT data to analyse their properties.  In the G-band, orphan and regular penumbrae have the same morphology, although the former are slightly brighter on average. The intensity difference is more pronounced in the upper photosphere, which is probed by the \\ion{Ca}{ii}~H filtergrams (Figs.~\\ref{histograms}a and~\\ref{histograms}c).  We suspect the reason for the enhanced intensity of the filaments in orphan penumbrae is their more elevated geometrical height, as they are not limited by the magnetic canopy of the background field present in sunspots.  Although there are some differences in the amplitudes of apparent motions in orphan and regular penumbrae, their overall structure is the same in the low and upper photosphere.  We observe apparent motions opposite to the gas flow in the inner penumbra (heads of the filaments) and apparent motions following the flow in the outer parts of orphan and regular penumbrae (Figs.~\\ref{histograms}b and~\\ref{histograms}d). Similarly to sunspots, there is a moat region near orphan penumbrae, which is located in the area towards which the flow is directed (extending the relation between the Eveshed flow and moat region found by \\citealt{Vargas:2007} to orphan penumbrae).  The analysis of Hinode SP and NFI data shows that orphan penumbrae are magnetic structures with the shape of flat $\\Omega$~loops. During the evolution of orphan penumbrae, the opposite polarities flanking them approach each other, while the filaments get shorter and disappear, which points to the submergence of the whole structure. Our velocity measurements cannot be used to demonstrate the downward motion as they are dominated by ``Evershed'' flows, and even observations taken close to the disc centre do not exhibit downward motions. We also cannot distinguish whether the filaments disappear earlier in the \\ion{Ca}{II}~H images or the G-band filtergrams.  The LOS velocities observed in orphan penumbrae are larger than those of nearby regular penumbrae, but the structure of the flow is the same. The lowest velocities occur near the filament heads, while the flow speed increases along the filaments and reaches a maximum in the tails. Using full Stokes polarimetry, we were able to determine the orientation and strength of the magnetic field along individual filaments of orphan penumbrae (Figs.~\\ref{detail_op}, \\ref{sp_evol}, \\ref{siphon_map}, and~\\ref{siphon_plot}). We found that the flows are oriented towards the footpoints with a stronger magnetic field, which agrees with the siphon flow mechanism as the driver of the flow.  Due to the similarities between orphan and regular penumbrae, we tentatively propose that orphan penumbrae represent the weaker and more inclined magnetic field component of regular penumbrae \\citep[not influenced by the background field present in sunspots,][]{Solanki:1993}.  On the other hand, this paper supports the siphon flow mechanism as the driver of the Evershed flow.  It remains to be studied if the observed magnetic field configuration is just a consequence of the fast flow \\citep[as suggested by][]{Rempel:2011} or the actual driver of the flow."
    },
    "1402/1402.4389_arXiv.txt": {
        "abstract": "A small blowout jet was observed at the boundary of the south polar coronal hole on 2011 February 8 at around 21:00~UT. Images from the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamics Observatory (SDO) revealed an expanding loop rising from one footpoint of a compact, bipolar bright point. Magnetograms from the Helioseismic Magnetic Imager (HMI) on board SDO showed that the jet was triggered by the cancelation of a parasitic positive polarity feature near the negative pole of the bright point. The jet emission was present for 25~mins and it extended 30~Mm from the bright point. Spectra from the EUV Imaging Spectrometer on board \\hinode\\ yielded a temperature and density of 1.6~MK and 0.9--1.7 $\\times$ $10^8$~cm$^{-3}$ for the ejected plasma. Line-of-sight velocities reached up to 250~\\kms\\ and were found to increase with height, suggesting plasma acceleration within the body of the jet. Evidence was found for twisting motions within the jet based on variations of the LOS velocities across the jet width. The derived angular speed was in the range 9--$12\\times 10^{-3}$~rad~s$^{-1}$, consistent with previous measurements from jets. The density of the bright point was 7.6 $\\times$ $10^8$~cm$^{-3}$, and the peak of the bright point's emission measure occurred at 1.3~MK, with no plasma above 3~MK. ",
        "introduction": "Coronal jets produce hot plasma, $>$~1~MK, that is ejected upwards away from the solar surface at high speeds up to several hundred \\kms, and they have been identified in coronal holes, quiet Sun and active regions \\citep{shimojo96}. The distinctive shape of jets is believed to be due to small-scale closed magnetic loops interacting with larger-scale open field structures \\citep{shibata86} and, as such, they are important as a means for studying basic plasma heating and acceleration mechanisms. All jets have their origins in a bright point on the solar disk. Often the bright point will have a simple, bipolar morphology (particularly in the case of coronal hole jets) and the coronal loops of the bright point are referred to as the base arch. The name `jet' is sometimes used to refer to the entirety of the event, with the term `spire' used to refer to the plasma ejected upwards from the bright point.  Recently \\citet{moore10} suggested the existence of two types of coronal hole jet, referred to as \\emph{standard} and \\emph{blowout}. The distinction was further clarified in \\citet{moore13}, who identified standard jets as having a single, narrow spire and a brightening at the edge of the bright point. Blowout jets have more complex spires (often a `curtain' of emission is seen) and the interior of the bright point intensifies. Subsequently, a number of authors have attempted to interpret their jet observations (both inside and outside of coronal holes) in terms of this dichotomy and it is apparent that a number of jets show complex behavior not easily assigned to either group, particularly when the high spatial and temporal resolution data from the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamics Observatory (SDO) are used. The blowout jets have been interpreted as being small-scale coronal mass ejections \\citep[micro-, or mini-CMEs,][]{nistico09,hong11}, with cool plasma from a mini-filament ejected with the coronal plasma, although this was not a defining characteristic for \\citet{moore13}, who classified events according to features seen in X-ray observations.   The launch of SDO in 2010 has been a boon for studies of transient events such as jets, as it yields full-disk solar images in multiple wavelengths at regular time cadences and with high spatial resolution. This ensures that any event on the Sun is almost certain to be captured with a consistent, high quality data-set. Jets seen with AIA often show complex ejecta \\citep{liu11, shen11, chen12, shen12, kslee13, schmieder13} and categorization into standard or blowout may not be straightforward \\citep{liu11,schmieder13}.  An important measurement that has been made from some SDO events is a twisting motion of the spire, identified as motions of compact bright features within the ejecta that can be tracked to reveal motions transverse to the jet axis \\citep{liu11,shen11,chen12,hong13}. Twist is an important parameter for understanding the physical mechanism behind the jet evolution as it can arise from emerging flux \\citep{pariat09}, and \\citet{hong13} have suggested that twist in blowout jets could arise from the erupting mini-filament associated with the event.   Most of the jets studied with SDO data have been large events seen in the quiet Sun or near active regions. Coronal holes hold a particular advantage for studying jets as the coronal emission above 1~MK is weak, and so ejecta that are hotter than this can be observed clearly. Additionally the magnetic structure in the photosphere is generally simpler, with the magnetic poles of the bright point lying within the unipolar magnetic field of the coronal hole. \\citet{shen11} and \\citet{hong13} reported jets seen in polar coronal holes, but in each case the jet was at the limb and so magnetic field data were unavailable. \\citet{young-blowout} presented SDO observations of a blowout jet occurring in a coronal hole on the disk, combined with EUV spectroscopic measurements from the EUV Imaging Spectrometer on board \\hinode. The jet was triggered by the two dominant polarities of the bright point converging and canceling with each other. The jet's spire was a broad curtain that extended over 70~Mm, and intense, small-scale kernels were seen within the bright point as the jet evolved. The spectroscopic data enabled the temperature and density of the jet plasma to be estimated at 1.4~MK and $2.7\\times 10^8$~cm$^{-3}$, respectively. The LOS jet speed was shown to reach up to 250~\\kms\\ and was found to increase with height.   The present work is a follow-up to that of \\citet{young-blowout}, using data from the same observation set but for a different event occurring at the coronal hole boundary. This event showed similarly large LOS velocities, but it was much smaller, it showed a different magnetic field evolution, and the morphology of the ejected plasma was different. Parameters of the jet measured from the AIA and Helioseismic Magnetic Imager (HMI) onboard SDO  and \\hinode/EIS are presented in Sects.~\\ref{sect.sdo} and \\ref{sect.eis}. A summary is given in Sect.\\ref{sect.summary} with a particular focus on comparisons with the jet of \\citet{young-blowout}. ",
        "conclusions": "\\label{sect.summary}  The key properties of the 8-Feb jet are as follows:  \\begin{itemize} \\item The jet is an expanding, loop-shaped structure that emerges from the negative polarity end of a small magnetic bipole. \\item The jet is triggered by cancellation of a weak, positive polarity magnetic feature that moves towards the negative pole of the bipole. \\item The bright point does not get hotter than 2.0~MK during the event, and the coronal density is $7.6\\times 10^8$~cm$^{-3}$. \\item The jet duration is 25~minutes, and the bright point is not significantly disrupted by the jet occurrence (the intensity returns to the pre-jet levels, and the magnetic field strength remains similar). \\item The jet extends 30~Mm (projected distance) above the bright point and the temperature of the ejected plasma is 1.7~MK; the density is in the range 0.9 to 1.7 $\\times$ $10^{8}$~cm$^{-3}$; and  absorption in the \\aix\\ channel suggests that chromospheric plasma is also ejected. \\item LOS speeds increase with distance from the bright point, reaching values up to  250~\\kms. \\item Evidence from the spectroscopic LOS velocities is found for twisted motions in the jet, and an angular speed of $9-12 \\times 10^{-3}$~rad~s$^{-1}$ is derived. \\end{itemize}    Comparing with the 9-Feb jet analyzed by \\citet{young-blowout}, the key difference lies with the magnetic field evolution: for the 9-Feb jet the two dominant polarities canceled with each other, whereas for the 8-Feb jet a parasitic polarity emerges near the negative polarity footpoint, and cancels with it. This difference is reflected in the evolution of the bright point during the event: the base arch of the 9-Feb jet is blown open by the jet, and a number of very intense kernels are seen; yet for the 8-Feb jet the base arch remains and no kernels were seen. Despite these differences, the temperature, density and velocity of the ejected plasma are similar between the two events.  On-disk observations of coronal holes afford an excellent opportunity for comparing different types of coronal jet, as both the  magnetic and coronal evolution can be studied. The results here and in \\citet{young-blowout} suggest that blowout jets can behave quite differently depending on the particular magnetic field geometry and evolution. Further categorization perhaps based on the types of magnetic field interaction (majority--majority or minority--majority cancelation) may be appropriate.   \\bigskip  The authors acknowledge funding from National Science Foundation grant AGS-1159353, and we thank the referee for valuable comments. P.R.Y. thanks ISSI for financial support to attend the 2014 International Team Meeting ``Understanding Solar Jets and their Role in Atmospheric Structure and Dynamics'' (PI: N.-E.~Raouafi), and he thanks the participants for useful discussions. \\hinode\\ is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). SDO is a mission for NASA's Living With a Star program, and data are provided courtesy of NASA/SDO and the AIA and HMI science teams."
    },
    "1402/1402.5346_arXiv.txt": {
        "abstract": "We present a spectral decomposition technique and its applications to a sample of galaxies hosting large-scale counter-rotating stellar disks. Our spectral decomposition technique allows to separate and measure the kinematics and the properties of the stellar populations of both the two counter-rotating disks in the observed galaxies at the same time. Our results provide new insights on the epoch and mechanism of formation of these galaxies. ",
        "introduction": "Some galaxies, despite of their simple morphologies, can reveal to be very complex sistems from a kinematic point of view. Clear examples are objects with stars and/or gas in orthogonal rotation with respect the rest of the galaxy (e.g., \\citealt{Corsini+03, Coccato+04, Corsini+12}), or with a warped inner structure (e.g., \\citealt{Coccato+07}).  Among these galaxies with multi-spin components, {\\em counter-rotating galaxies\\/} are those that host two components that rotate in opposite directions from each other. These peculiar objects have been observed from in all morphological types. They are classed by the nature (stars vs. stars, stars vs. gas, gas vs. gas) and size (counter-ro\\-ta\\-ting cores, rings, and disks) of the counter-rotating components (see \\citealt{Bertola+99}, for a review).  The subject of this work is the peculiar case of galaxies hosting two counter-rotating stellar disks of comparable size.  The prototype of this class of objects is the famous E7/S0 galaxy NGC~4550, whose counter-rotating nature was first discovered by \\citet{Rubin+92}.  The study of the kinematics and the stellar population of the two counter-rotating components are the key to understand the formation mechanisms of these objects. This task is complicated by the fact that the two components are co-spatial, and therefore we observe their combined contribution at each position on the sky. Thus, it is necessary to separate their contribution to the total observed galaxy spectrum and measure the properties and kinematics of their stellar populations separately. To this aim we developed a spectroscopic decomposition technique and successfully applied to three spiral galaxies with counter-rotating stellar disks \\citep{Coccato+11, Coccato+13b, Coccato+13}.   \\begin{figure}[t!] \\centering \\includegraphics[width=13cm]{coccato_01_f01.eps} \\caption{Results of the spectroscopic decomposition in NGC~3593. Top panels: Observed field of view, absorption line indices, and prediction from stellar population models \\citep{Thomas+11}; red and blue symbols refer to the main and secondary stellar components, respectively. Bottom panels: 2-dimensional velocity fields of the main and secondary stellar components, and of the ionized gas component.} \\label{fig:003_1} \\end{figure} ",
        "conclusions": "Our results rule out internal processes, such as bar dissolution, since they would lead to the formation of two counter-rotating disks with the same properties.  External acquisition of gas followed by star formation, or mergers of two galaxies with similar mass are both viable mechanisms. The accretion scenario predicts that the secondary component is the younger, and that the two components have different metallicity, as they are formed from different gas clouds. In the second scenario, the properties of the two stellar component depend on the properties of the merging systems and star formation episodes triggered by the merger.  A larger sample is therefore required to statistically determine which mechanism is the most efficient to build counter-rotating stellar disks."
    },
    "1402/1402.3176_arXiv.txt": {
        "abstract": "{We investigate the probability of obtaining an observable curvature perturbation, using as an example the minimal curvaton-higgs (MCH) model. We determine ``probably observable'' and ``probably excluded'' regions of parameter space assuming generic initial conditions and applying a stochastic approach for the curvaton's evolution during inflation. Inflation is assumed to last longer than the $N_{obs} \\simeq 55$ observable $e$-folds, and the {\\em total} number of $e$-folds of inflation determines the particular ranges of parameters that are probable. For the MCH model, these ``probably observable'' regions always lie within the range $8\\times 10^4 \\GeV \\leq m_\\sigma \\leq 2\\times 10^7\\GeV$, where $m_\\sigma$ is the curvaton mass, and the Hubble scale at horizon exit is chosen as $H_* = 10^{10}\\GeV$. Because the ``probably observable'' region depends on the total duration of inflation, information on parameters in the Lagrangian from particle physics and from precision CMB observations can therefore provide information about the total duration of inflation, not just the last $N_{obs}$ $e$-folds. This method could also be applied to any model that contains additional scalar fields to determine the probability that these scalar fields contribute to the curvature perturbation.}  \\begin{document} ",
        "introduction": "\\label{sec:intro} In order to generate sufficient inflation, and in order to obtain the observed amplitude of scalar perturbations, an inflationary theory must give certain initial conditions to the fields involved. In this paper we focus on this problem as applied only to spectator fields, and study the initial field value of the curvaton field.  The curvaton scenario is an alternative mechanism to generate the primordial curvature perturbation $\\zeta$ \\cite{curvorig}. It is interesting both because it can allow many inflaton models which are successful for generating inflation but cannot produce sufficient perturbations, and because curvaton models can produce features such as non-Gaussianity and isocurvature, allowing constraints from data to be placed. In addition, the coupling of the curvaton to other fields is less tightly constrained than for the inflaton, allowing for stronger couplings to the standard model.  Although the curvaton model has been well-studied, it is often suggested that the scenario is fine-tuned or unnatural. The parameters of a typical curvaton model are: the Hubble parameter $H_*$, which gives the energy scale of inflation; the effective curvaton decay width $\\Gamma$; the parameters of the potential such as the mass $m_\\sigma$ and quartic self-coupling $\\lambda_\\sigma$; and the value of the homogeneous curvaton field during inflation $\\sigma_*$, determined as observable scales exit the horizon approximately $N_{obs}\\simeq55$ $e$-folds before the end of inflation. Generally, by tuning either $\\Gamma$ or $\\sigma_*$, the correct amplitude of curvature perturbation can be obtained, $\\zeta =4.7\\times 10^{-5}$ \\cite{Planck:zeta}. Some attempts have been made to determine $\\Gamma$ in terms of the particle physics couplings of the model. A specific example is the non-perturbative decay of the curvaton into standard model Higgs bosons \\cite{Enqvist:2012tc,Enqvist:2013qba,Enqvist:2013gwf,Mukaida:2014yia}; the perturbative case is discussed in \\cite{Enqvist:2011jf}.     In this paper, we investigate the probability of obtaining the measured curvature perturbation in the minimal curvaton-higgs (MCH) model. In this model, the curvaton is a real singlet scalar, coupled to the standard model higgs with the Lagrangian in the unitary gauge given by \\be{Lagrangian} {\\cal L} = {\\cal L}_{SM} + \\hf \\partial_\\mu \\sigma \\partial^\\mu \\sigma +  \\frac{1}{2}m_\\sigma^2 \\sigma^2 + \\frac{\\lambdasigma}{4!} \\sigma^4 + \\frac{1}{4} g^2\\sigma^2h^2, \\ee where ${\\cal L}_{SM}$ is the standard model Lagrangian including the higgs and $h$ is the physical higgs boson. Because the higgs has a large field value during inflation, $h \\gg 246\\GeV$,  the higgs potential is well approximated by \\be{higgs_potential} V(h) = \\frac{\\lambda(\\mu)}{4} h^4, \\ee where $\\lambda(\\mu)$ is the higgs self coupling at the relevant energy scale. We now assume that the curvaton's effective quartic coupling is negligible, i.e.\\ $\\lambdasigma = 0$.  In general, effective quartic terms are generated by loop corrections\\footnote{See \\cite{Enqvist:2011jf} for a discussion of this in the curvaton scenario.}, however, setting $\\lambdasigma = 0$ enables us to calculate the stochastic behaviour (mostly) analytically.   We assume that inflation lasts for a total duration of $N = N_{pre} + N_{obs}$ $e$-folds. We then use a stochastic approach \\cite{Starobinsky86} to quantify the probability of obtaining particular values of $\\sigma_*$, calculated $N_{obs}$ $e$-folds before inflation ends. Because the observables $\\zeta$ (amplitude of the curvature perturbation) and $\\fnl$ (non-Gaussianity) depend on $\\sigma_*$, we can then translate the probability distribution of $\\sigma_*$ into a probability for obtaining any particular range of these observables. This particular approach was first suggested in \\cite{Enqvist:2012xn}, although applying the equilibrium limit of the stochastic dynamics to the curvaton model was already considered in \\cite{Lyth:FP}. The influence of superhorizon perturbations on observables has also been discussed, e.g.\\ in \\cite{superhor}.  We find that for most values of $N_{pre}$ there are wide ranges of $g$ and $m_\\sigma$ that give a high ($>10\\%$) probability of producing a curvature perturbation that is \\obs,\\ i.e.\\ neither too large to be excluded, nor too small to be negligible. This is significant because it enables us to relate the parameters of the model to knowledge about the duration of inflation, which could be found from a fundamental theory. Conversely, information from particle physics and CMB observations could give us information about the duration of inflation before the last $N_{obs}$ $e$-folds, via their predictions for curvaton model parameters.  This methodology could also be used to quantify, under various assumptions, whether additional light scalars in any particular model would be likely to give a non-negligible contribution to predictions.  The paper is organised as follows. In \\sect{sec:dyn} we introduce the relevant stochastic evolution equation, discuss the decay mechanisms for the curvaton field, and calculate the contribution of the higgs field during inflation to the curvaton's effective mass. In \\sect{sec:prob} we calculate the probability that the model gives $\\zeta$ and $\\fnl$ that are either \\obs,\\ \\negl\\ or \\excl\\ by current observations. Then in \\sect{sec:results} we present our results first for the case $N_{pre} \\to \\infty$ and secondly for the case of finite $N_{pre}$. Finally in \\sect{sec:conc} we conclude. ",
        "conclusions": "\\label{sec:conc}  We can think about spectator fields during inflation in two ways. The first is that if (non-inflaton) scalar fields exist in a theory, do they either rule out the theory or otherwise affect observational predictions? The second is that if we design a model such that a spectator field is responsible for the observed density perturbation, does this occur with natural or with fine-tuned initial conditions? This paper provides a framework to answer both of these questions. In the absence of fine-tuned initial conditions, we find in almost all cases that probable regions of parameter space exist for some range of model parameters (considering the minimal curvaton-higgs (MCH) model \\cite{Enqvist:2013gwf} with $H_* = 10^{10}\\GeV$). In addition, large masses $m_\\sigma\\gtrsim 2\\times 10^7\\GeV$ are disfavoured due to negligible curvature perturbations.  We obtained our results by assuming the curvaton had the initial condition $\\sigma_0 = 0$ (\\fig{fig:Nsig_0}) or $\\sigma_0 = \\mpl$ (\\fig{fig:N_sigMpl}) before inflation. We assumed that inflation lasted $N = N_{pre} + N_{obs}$ $e$-folds and considered various values of $N_{pre}$ including the $N_{pre}\\to \\infty$ limit (\\fig{fig:higgscont}). From the initial strongly peaked $\\sigma_0$, the long wavelength modes of the curvaton field behave stochastically. We calculated the probability distribution for the mean field value $\\sigma$ at a given time; the observable Universe is given by one particular realisation of this distribution. We then translated this probability distribution to a probability of obtaining particular ranges for the curvature perturbation $\\zeta$ and the non-Gaussianity $\\fnl$. In this way we attempted to quantify how likely a set of parameters are to give an \\obs,\\ \\negl\\ or \\excl\\ Universe.  We used the MCH model \\cite{Enqvist:2013gwf} throughout the paper to illustrate our method. In this model the curvaton has a quadratic potential and is coupled only to the standard model Higgs Boson with coupling $g$. For simplicity, we ignored a number of effects, including the effective quartic coupling of the curvaton and the possibility of an inflaton-curvaton coupling. We also assumed that the dominant mechanism for curvaton decay was through dimension-5 non-renormalisable couplings. At $g\\lesssim (m_\\sigma/\\mpl)^{1/4}$ this assumption is likely to be reasonable, but the behaviour of such an oscillating field in a thermal background after inflation is complicated and is the subject of ongoing study  (see \\cite{D'Onofrio:2012qy,Enqvist:2013qba,BasteroGil,Drewes:2013iaa,Mukaida:2014yia,Mukaida,Enqvist:2012tc}). For example, the generation of perturbations in a case similar to the MCH model could also involve ``self-modulated preheating'' \\cite{Mukaida:2014yia}. Including all of these effects in our calculations would certainly be interesting, although it would not be possible within our semi-analytical framework.  We found that $m_\\sigma > 2\\times 10^{7} \\GeV$ always gives a Universe with negligible $\\zeta$ with $>99\\%$ probability. Masses below $m_\\sigma = 8\\times 10^4\\GeV$ are excluded due to a late curvaton decay spoiling the predictions of BBN. Considering $8\\times 10^{4}\\GeV \\leq m_\\sigma \\leq 2\\times 10^7\\GeV$, observable $\\zeta$ has high probability ($> 10\\%$ and even $> 50\\%$) for broad ranges of $m_\\sigma$ and $g$. A notable region with small $m_\\sigma$ and large $g$  is likely to give a too large curvature perturbation. There is very little dependence on the initial conditions for large effective mass $m_{eff}^2 = m_\\sigma^2 + g^2 h_*^2$ because the equilibrium distribution is already achieved for small $\\Npre$. For smaller $m_{eff}$ (small $m_\\sigma$ and $g$), there is sensitivity to the initial field value $\\sigma_0$ and to the duration of inflation. In general, larger values of $m_\\sigma$ become probable only after a longer duration of inflation. All of these conclusions are for $H_* = 10^{10}\\GeV$.      Future experiments, observations and theoretical progress are expected to further constrain the MCH model. For example, the Planck satellite will be sensitive to non-Gaussianity down to $\\fnl \\simeq 5$ and can also constrain $H_*$. Further theoretical calculations of curvaton decay mechanisms are likely to constrain $m_\\sigma$ and $g$. After an inflation model is assumed, $m_\\sigma$ and $g$ may also be constrained by measurements of the CMB spectral index.  Additionally, if WIMP dark matter were to be favoured, the curvaton would be required to decay before dark matter freeze-out to avoid large isocurvature --- and because the dimension-5 decay width depends on $m_\\sigma$, a large part of the low $m_\\sigma$ parameter space would then be ruled out. Together, these advances may make it possible to link the MCH model with the duration of inflation. For example if the above considerations favoured large $m_\\sigma$, a curvaton model without fine-tuning would require a long duration of inflation. If small $m_\\sigma$ and large $g$ was instead favoured, $N_{pre} = {\\cal O}(10-100)$ would be sufficient for a ``probable'' Universe, provided that $\\sigma_0 = 0$ could be justified. Conversely, knowledge from fundamental theories about the total duration of inflation could also give information about the parameters in the curvaton model.    In summary, we have studied the ``naturalness'' of the MCH scenario under generic assumptions about the curvaton initial conditions and various other simplifications. We find that for $8\\times 10^4\\GeV \\leq m_\\sigma \\leq 2\\times 10^7\\GeV$, the curvaton mechanism can ``naturally'' occur, provided that inflation has lasted long enough. Our methodology could be applied to more general curvaton models, as well as to models with additional spectator fields. It could then be determined whether these additional fields must have fine-tuned initial conditions to avoid spoiling the predictions of the model."
    },
    "1402/1402.1173_arXiv.txt": {
        "abstract": "We consider an effective field theory for a gauge singlet Dirac dark matter (DM) particle interacting with the Standard Model (SM) fields via effective operators suppressed by the scale $\\Lambda \\gtrsim 1$~TeV. We perform a systematic analysis of the leading loop contributions to spin-independent (SI) DM--nucleon scattering using renormalization group evolution between $\\Lambda$ and the low-energy scale probed by direct detection experiments. We find that electroweak interactions induce operator mixings such that operators that are naively velocity-suppressed and spin-dependent can actually contribute to SI scattering. This allows us to put novel constraints on Wilson coefficients that were so far poorly bounded by direct detection. Constraints from current searches are already significantly stronger than LHC bounds, and will improve in the near future. Interestingly, the loop contribution we find is isospin violating even if the underlying theory is isospin conserving. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1345_arXiv.txt": {
        "abstract": "The existence of the Standard Model (SM) Higgs implies that a gravitational wave (GW) background is generated by the decay products of the Higgs, soon after the end of inflation. % Theoretically, all Yukawa and $SU(2)_{\\rm L}$ gauge couplings of the SM are imprinted as features in the GW spectrum. In practice, the signal from the most strongly coupled species dominates, rendering inaccesible the information on the other species. If detected, this background could be used for measuring properties of high-energy particle physics, including beyond the SM. To achieve this goal, new high frequency GW detection technology is required, beyond that of currently planned detectors. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0343_arXiv.txt": {
        "abstract": "Planetary and stellar magnetic fields are thought to be sustained by helical motions ($\\alpha$-effect) and, if present, differential rotation ($\\Omega$-effect). In the Sun, the strong differential rotation in the tachocline is responsible for an efficient $\\Omega$-effect creating a strong axisymmetric azimuthal magnetic field. This is a prerequisite for Parker dynamo waves that may be responsible for the solar cycle. In the liquid iron cores of terrestrial planets, the Coriolis force organizes convection into columns with a strong helical flow component. These likely dominate magnetic field generation while the $\\Omega$-effect is of secondary importance. Here we use numerical simulations to show that the planetary dynamo scenario may change when the heat flux through the outer boundary is higher in one hemisphere than in the other. A hemispherical dynamo is promoted that is dominated by fierce thermal wind responsible for a strong $\\Omega$-effect. As a consequence Parker dynamo waves are excited equivalent to those predicted for the Sun. They obey the same dispersion relation and propagation characteristics. We suggest that Parker waves may therefore also play a role in planetary dynamos for all scenarios where zonal flows become an important part of convective motions. ",
        "introduction": "The dynamo mechanism of stars and planets relies on electromagnetic induction and requires an electrically conducting medium and a complex fluid flow to maintain the magnetic fields against Ohmic decay. The liquid iron cores form the dynamo regions of terrestrial planets and the dominance of Coriolis forces guarantees sufficiently complicated convective dynamics in these fast rotating bodies. The flow is organized in convective columns parallel to the rotation axis. Secondary flows up and down these columns yield a strong helical component essential for magnetic field generation. Since this helicity associated with the local columnar motion dominates the production of the poloidal and the toroidal magnetic field contributions in many dynamo simulations they are typically classified as $\\alpha^2$ dynamos \\cite{Olson1999}. The terminology goes back to the mean-field dynamo theory where an $\\alpha$ stands for the creation of large-scale field by small-scale motions.  In the Sun, the differential rotation at the tachocline is the main source of large-scale axisymmetric toroidal magnetic field. Since this is called an $\\Omega$-effect in mean-field dynamo theory, the solar dynamo is classified as $\\alpha \\Omega$ \\cite{bKrause1980}. A powerful $\\Omega$-effect may be the reason for the oscillatory nature of the solar magnetic field which has not been observed in planetary dynamos. Strong axisymmetric toroidal fields are a prerequisite for Parker dynamo waves  \\cite{Parker1955} which seem to mimic the general oscillating behaviour to a large extent.  In the framework of the classical mean-field model of an axisymmetric $\\alpha \\Omega$-dynamo \\cite{Parker1955}, Parker waves are purely magnetic oscillations between the poloidal and toroidal mean magnetic field components. Poloidal field is exclusively created by the $\\alpha$-effect as the parametrization of the nonaxisymmetric helical flows, whereas the azimuthal toroidal field is induced by the axisymmetric differential rotation or the $\\Omega$-effect \\cite{Tobias2007}.  The minimum frequency of these waves is simply the inverse of the magnetic diffusion time $\\tau_\\lambda=D^2/\\lambda$ \\cite{bruediger}. Assuming a turbulent bulk magnetic diffusivity of $\\lambda=10^{8}\\, \\mathrm{m^2/s}$ for the solar convective shell of thickness $D=2 \\times 10^{8} \\, \\mathrm{m}$, the magnetic diffusion time is $\\tau_\\lambda=12.7\\, \\mathrm{yr}$ and matches the time scale of a solar cycle ($22\\, \\mathrm{yr}$) reasonably well.  Parker waves are typically associated with dynamo models where stress free flow boundary conditions allow strong zonal winds to arise. These guarantee a strong $\\Omega$-effect and thereby promote Parker waves. Parker-wave like oscillation in combination with a hemispherical magnetic field have previously been reported by \\cite{Grote2000}. Because of the stress free flow boundary conditions and the fixed temperature conditions their model also seems more appropriate for a stellar application.  Here we report that Parker waves also appear in simulations of planetary dynamos when lateral variations in the heat flux through the outer boundary lead to fierce zonal winds and thus strong $\\Omega$-effects. More specifically, we consider a model for the ceased ancient Martian dynamo which left its trace in form of a remanent crustal magnetization. The magnetization shows a hemispherical distribution with much higher values in the southern than in the northern hemisphere \\cite{Acuna1999}. Recent studies aim at explaining this pattern with the fact that the dynamo itself produced a hemispherical field \\cite{Stanley2008,Dietrich2013}. Such dynamos are promoted when the heat flux through the southern core-mantle boundary is significantly larger than that through the northern. In terrestrial planets, the core-mantle-boundary heat flux is determined by the thermal mantle structure which develops on much slower time scales than core convection. A major impact in the northern hemisphere or a large-scale mantle convection pattern are two alternative scenarios causing a north/south heat flux asymmetry \\cite{Stanley2008}.  Since the southern hemisphere of the core is cooled more efficiently than its northern counterpart, a latitudinal temperature gradient develops that drives fierce thermal winds.  When the heat flux asymmetry is large enough these winds clearly dominate the flow. The dynamo then changes from an $\\alpha^2$ to an $\\alpha\\Omega$ type and shows wave-like behaviour reminiscent of Parker waves \\cite{Dietrich2013}. Here we analyse a suite of hemispherical dynamo simulations to determine when these waves arise. By showing that the frequency of the waves follows the theoretical dispersion relation based on mean-field theory we clearly establish their Parker wave nature. ",
        "conclusions": "We demonstrated that the oscillating solution found in strongly hemispherical dynamo solutions promoted by a north/south heat flux asymmetry through the outer boundary are likely Parker dynamo waves. The heat flux pattern leaves the northern hemisphere hotter than the southern and the associated temperature gradient drives fierce zonal thermal winds. These in turn lead to a strong $\\Omega$-effect, a prerequisite for Parker waves.  In the Martian context this is bad news because the relatively fast oscillations on the order of a few thousand years are not compatible with the high crustal magnetization amplitude \\cite{Dietrich2013}. Since crustal magnetization is acquired over several million years the effective magnetization as observed from a space craft would be very weak.  Lateral heat flux variations with $q^\\star\\approx 1$ seem generally possible for terrestrial planets, at least as long as the mantle is actively convecting. The large-scale cosine-like variation explored here, however, is geared to explain the Martian hemisphericity. For Earth, the heat flux pattern is dominated by a spherical harmonic of degree and order two. This was likely different in the past, but a degree one pattern nevertheless seems more likely for Mars while higher harmonics are expected for Earth \\cite{Roberts2006}. Another property essential for promoting strong thermal winds is the dynamo heating mode. The effect of an outer boundary heat flux variation is much larger for a dynamo without an inner core that is exclusively driven by secular cooling \\cite{Hori2012}. Whether a combination of a more complex heat flux pattern and secular cooling would also lead to strong thermal winds and possibly Parker waves has not been explored in any detail so far. This may have been the case for the geodynamo before the onset of inner core growth which may have happened only one Gyr ago.  Recent simulations geared to model the dynamo in the gas giants also show oscillatory behaviour reminiscent of Parker waves \\cite{Gastine2012}. Here the stress free outer boundaries allow much stronger zonal winds to develop than in typical simulations for terrestrial dynamos. This can also lead to a strong $\\Omega$-effect and associated Parker wave behaviour. Parker waves are thus not only interesting to explain the solar cycle but also have to be considered in the planetary context."
    },
    "1402/1402.0669_arXiv.txt": {
        "abstract": "{ The early evolution of a dense young star cluster (YSC) depends on the intricate connection between stellar evolution and dynamical processes. Thus, N-body simulations of YSCs must account for both aspects. We discuss N-body simulations of YSCs with three different metallicities ($Z=0.01$, 0.1 and 1 Z$_\\odot$), including metallicity-dependent stellar evolution recipes and metallicity-dependent prescriptions for stellar winds and remnant formation.  We show that mass-loss by stellar winds influences the reversal of core collapse. % In particular, the post-collapse expansion of the core is faster in metal-rich YSCs than in metal-poor YSCs, because the former lose more mass (through stellar winds) than the latter. As a consequence, the half-mass radius expands more in metal-poor YSCs. % We also discuss how these findings depend on the total mass and on the virial radius of the YSC. These results give us a clue to understand the early evolution of YSCs with different metallicity. } ",
        "introduction": "\\label{sec:Introduction} The densest young star clusters (YSCs) are collisional environments: their central two-body relaxation timescale ($t_{\\rm rlx}$) is generally shorter than $100$ Myr. For a realistic initial mass function (IMF), YSCs are expected to undergo core collapse (CC) in less than $t_{\\rm rlx}$. % The reversal of CC is generally ascribed to hard binaries (i.e. binaries whose binding energy is higher than the average kinetic energy of a star in the cluster), because they transfer kinetic energy to stars through three-body encounters.   It has long been debated whether mass-loss by stellar winds and/or supernovae (SNe) is efficient in affecting CC (e.g. \\citealt{mapelli&bressan2013}, and references therein). Stellar winds and SNe eject mass from a star cluster, making the central potential well shallower and quenching the onset of gravothermal instability. Furthermore, stellar winds are expected to depend on  metallicity ($Z$): metal-poor stars lose less mass than metal-rich ones (e.g. \\citealt{vink2001}). Thus, the effect of stellar winds on CC is expected to be stronger at high $Z$. In this paper, we investigate the impact of $Z$-dependent stellar winds and SNe on the CC of YSCs, by means of direct-summation N-body simulations. ",
        "conclusions": "We have shown that stellar winds are very important in the early evolution of YSCs and that their effects strongly depend on $Z$. In particular, the post-collapse re-expansion of the core is faster for metal-rich YSCs than for metal-poor YSCs, because the former lose more mass (through stellar winds) than the latter. As a consequence, the half-mass radius and the half-light radius expand faster in metal-poor YSCs. The initial size of the YSC plays a critical role, because the relaxation timescale and thus the onset of CC strongly depend on it ($t_{\\rm rlx}\\propto{}r_{\\rm hm}^{3/2}$). The total mass of the YSC has only marginal effects. Other YSC properties (e.g. $W_0$) deserve further investigation. Furthermore, the $Z$-dependence of stellar winds is still barely understood, especially in the post-main sequence evolution. Thus, forthcoming studies will also investigate the effects of different recipes of mass-loss by stellar winds.   \\begin{acknowledgement} MM acknowledges financial support from MIUR through grant FIRB 2012 RBFR12PM1F, from INAF through grant PRIN-2011-1 and from CONACyT through grant 169554. The authors thank the organizers and the participants of the fifth Guillermo Haro conference for the excellent organization and for the stimulating discussions. \\end{acknowledgement}"
    },
    "1402/1402.2493_arXiv.txt": {
        "abstract": "The analyses of Titan's gravity field obtained by Cassini space  mission suggest the presence of an internal ocean beneath its  icy surface. The characterization of the geophysical parameters of the icy shell and the ocean is important to constrain the evolution models of  Titan. The knowledge of the librations, that are periodic oscillations  around a uniform rotational motion, can bring piece of information on  the interior parameters.  The objective of this paper is to study the librational response in  longitude from an analytical approach for Titan composed of a deep atmosphere, an elastic icy shell, an internal  ocean, and an elastic rocky core perturbed  by the gravitational interactions with Saturn. We start from the librational equations developed for a rigid satellite in  synchronous spin-orbit resonance. We introduce explicitly the atmospheric torque acting on the surface computed from the Titan IPSL GCM \\textit{(Institut Pierre Simon Laplace  General Circulation Model)} and the periodic deformations of elastic solid layers due to the tides. We investigate the librational response for various interior models in order to compare and to identify the  influence of the geophysical parameters and the impact of the elasticity.  The main librations arise at two well-separated forcing frequency ranges: low forcing frequencies dominated by the Saturnian annual and semi-annual frequencies, and a high forcing frequency regime dominated by Titan's orbital frequency around Saturn. At low forcing frequency, the librational response is dominated by the Saturnian gravitational torque and the atmospheric torque has a small effect. In addition, the libration amplitude in that case is almost equal to the magnitude of the perturbation. The modulation of the gravitational torque amplitude at the orbital frequency with periodic deformation induces long-period terms in the librational response which contain information on the internal structure. At high forcing frequency the libration depends on the inertia of the layers and the elasticity can strongly reduce its amplitude at orbital frequency. For example, the amplitude of diurnal libration for oceanic models goes from about $320-390$ meters if the icy shell is purely rigid to $60-85$ meters when the elasticity is included, \\textit{i.e.} a reduction of about $80\\%$. For models without ocean, diurnal libration goes from $52$ meters in a rigid case to $50$ meters for an elastic case, a very low reduction due to the weak deformation of an entirely solid satellite compared to the deformation of a thin icy shell. Oceanic models with elastic solid layers have the same order of libration amplitude than the oceanless models, which makes more challenging to differentiate them by the interpretation of librational motion. ",
        "introduction": "\\label{sec:mis}  Titan's internal structure has been revealed by accurately tracking the trajectory of Cassini spacecraft approaching Titan during six flybys \\citep{Iess10, Iess12}. They measured the gravity field and its variations allowing to infer information on the density profile of the satellite. Since the inversion between the gravity field and density profile is not unique, only a range of models can be determined. The add-on assumption that Titan is close to the hydrostatic equilibrium led \\cite{Iess10} to obtain an estimation of the moment of inertia (MoI) $I$ of Titan between $0.33$ and $0.34\\ MR^2$ (where $M$ and $R$ are the mass and mean radius of Titan, respectively). Such small MoI implies an increase of density towards Titan's center and requires a low-density core to match the mean density of $1881$ kg m$^{-3}$ \\citep{Iess10}.  In parallel to these gravity measurements, two categories of thermal and chemical models have been developed to determine Titan's internal structure (\\textit{e.g.} \\cite{Castillo10,Castillo12, McKinnon11, Fortes12} and see the review of \\cite{Tobie12}). In the first category, as developed by \\cite{Fortes12}, the models described essentially the upper layers (dense and light ocean, comparison of solid models with pure water ice or methane clathrate layers), while in the second category the models focused on the inner core composition assuming the presence of a global ocean layer (\\textit{e.g.} \\cite{Castillo12, McKinnon11}).  \\cite{Castillo12} built interior models made of anhydrous silicate core surrounded by hydrated rock and water ice while  \\cite{McKinnon11} focused on hydrated silicates surrounded by mixture of rock and ice. These models of inner core allow the presence of a small iron core, and are compatible with an internal ocean as deduced by \\cite{Iess12}.  Six different models reported in Table \\ref{table1} have been selected, coming from \\cite{Fortes12} (the models called F1, F2 and the solid model F3), \\cite{Castillo10} (models CA10 and FE10) and \\cite{McKinnon11} (model MC11). The FE10 model is similar to the CA10 model with an additional small inner iron core as suggested by \\cite{Castillo10}. The icy shell thickness has been taken equal to $100$ kilometers for each model. Such value corresponds to the upper bound obtained by topographical model of \\citep{Nimmo10}. The lower bound has been obtained by \\cite{Beghin12} by using the Schumann's resonance in Titan's atmosphere. The icy shell thickness is an essential parameter for Titan's models with a rigid shell \\citep{Vanhoolst09} but the elasticity strongly diminish its influence (see Section \\ref{sec:libration}). \\cite{Fortes12} used oceans with a bottom mean radius of $2225$ kilometers (models F1 and F2). To be able to compare the influence of the ocean density and the inner core structure on the libration, the same ocean bottom mean radius is used for the CA10, MC11 and FE10 models. For a given set of solid layers sizes and densities, the ocean density is then adjusted to conserve Titan's MoI and mass.   \\begin{table*}[] \\footnotesize \\caption{\\label{table1} Characteristics of the internal structure models selected for this paper. The first three models quoted \\textit{Light Ocean} (F1), \\textit{Dense Ocean} (F2) and \\textit{Pure-water-ice} (F3) models are from \\cite{Fortes12}. Model CA10 is from \\cite{Castillo10} and MC11 from \\cite{McKinnon11}, both with dense ocean and high pressure ice layers. The last model (FE10) is a model similar to CA10 including an iron core. $I$ is the total moment of inertia of the body, while $C_s$ and $C_i$ are polar moments of inertia of the icy shell and solid interior respectively, expressed in terms of mass and radius of Titan ($M_T=13455.3\\ 10^{19}$ kg and $R_T=2575.5$ km following \\cite{Fortes12}).} \\begin{center} \\begin{tabular}{l c c c c c c} \\hline & \\multicolumn{2}{c}{\\begin{bf} F1 \\end{bf}}  & \\multicolumn{2}{c}{\\begin{bf} F2 \\end{bf} } &  \\multicolumn{2}{c}{\\begin{bf} F3 \\end{bf}}\\cr \\hline Layer & $\\rho$ (kg m$^{-3}$)  & R (km) & $\\rho$ (kg m$^{-3}$) & R (km)& $\\rho$ (kg m$^{-3}$) & R (km)\\cr \\hline Icy shell & 930.9 & 2575.5 & 930.9 & 2575.5 &  932.8 & 2575.5 \\cr Ocean & 1023.5 & 2475 & 1281.3 & 2475 & - & - \\cr HP ice 1 & - & - & - & - & 1193.3 & 2429  \\cr HP ice mantle & 1272.7 & 2225 & 1350.9 & 2225 & 1268.7 & 2304  \\cr Silicate mantle & 1338.9 & 2163 & - & - &1338.9 & 2185 \\cr Silicate core & 2542.3 & 2116 & 2650.4 & 1984 & 2584.1 & 2055\\cr \\hline $I$ (MR$^2$) & \\multicolumn{2}{c}{0.3414}&\\multicolumn{2}{c}{0.3413}&\\multicolumn{2}{c}{0.3413}\\cr $C_s$ (MR$^2$) & \\multicolumn{2}{c}{0.0357}&\\multicolumn{2}{c}{0.0357}&\\multicolumn{2}{c}{-}\\cr $C_i$ (MR$^2$) & \\multicolumn{2}{c}{0.2320}&\\multicolumn{2}{c}{0.2133}&\\multicolumn{2}{c}{-}\\cr \\hline \\cr  & \\multicolumn{2}{c}{\\begin{bf} CA10 \\end{bf}}  & \\multicolumn{2}{c}{\\begin{bf} MC11 \\end{bf} } &  \\multicolumn{2}{c}{\\begin{bf} FE10 \\end{bf}}\\cr \\hline Layer & $\\rho$ (kg m$^{-3}$)  & R (km) & $\\rho$ (kg m$^{-3}$) & R (km)& $\\rho$ (kg m$^{-3}$) & R (km)\\cr \\hline Icy shell & 932.8 & 2575.5 & 932.8 & 2575.5 &  932.8 & 2575.5 \\cr Ocean & 1317.6 & 2475 & 1315.5 & 2475 & 1224.1 & 2475 \\cr HP ice mantle & 1350.9 & 2225 & 1350.9 & 2225 & 1350.9 & 2225  \\cr Rocky mantle & 2520 & 2000 & 2200 & 2075 & 2520 & 2000 \\cr Silicate core & 3400 & 900 & 3300 & 1300 & 3400 & 900\\cr Iron core & - & - & - & - & 6500 & 500 \\cr \\hline $I$ (MR$^2$) & \\multicolumn{2}{c}{0.3402}&\\multicolumn{2}{c}{0.3378}&\\multicolumn{2}{c}{0.3336}\\cr $C_s$ (MR$^2$) & \\multicolumn{2}{c}{0.0358}&\\multicolumn{2}{c}{0.0358}&\\multicolumn{2}{c}{0.0358}\\cr $C_i$ (MR$^2$)& \\multicolumn{2}{c}{0.2095}&\\multicolumn{2}{c}{0.2073}&\\multicolumn{2}{c}{0.2096}\\cr \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "In this paper, the librational response in longitude of Titan has been investigated by taking into account the presence of a global ocean, the atmospheric torque, the orbital perturbations and the elasticity of solid layers. A linearized model has been developped and the amplitudes of libration have been computed and decomposed into responses at different frequencies. The librational responses can be divided into three regimes depending on the forcing frequencies: the low forcing frequencies at Saturnian annual and semi-annual periods, the high forcing frequencies at Titan's orbital and semi-orbital periods, and the forcing frequency ranges close to the proper frequencies of the system.  For the forcing at low frequencies, the libration response of the satellite is dominated by the gravitational torque and the internal structure has a negligible effect. In the high frequency forcing case, the librational response depends on the inertia of the body. The rigid case presents a large librational response at orbital frequency that is strongly reduced by the introduction of elasticity as stated by \\cite{Goldreich10}. They introduced an elastic torque that corresponds to the pressure exerted by the ocean on the deformed icy shell as explained by \\cite{Goldreich10} in Appendix (A.6). In our approach, the expression of the pressure torque is obtained according to the procedure detailed in \\cite{Vanhoolst09}. The terms that result from the pressure torque are represented with a prime as shown in section 4.1 for the rigid case and section 4.2 for the purely non-rigid case. The pressure torque for the non-rigid case results from the deformation of the shell-ocean interface and it is characterized by the term $9/2 n^2 \\Delta I'_l$. Thus physically, the pressure torque $9/2 n^2 \\Delta I'_l$ and the pressure torque of \\cite{Goldreich10} have the same origin. Then, the literal comparison of the two expressions is difficult because the two approaches are different. \\cite{Goldreich10} assumed that the shell is a massless membrane overlaying a purely fluid body. In our case, Titan is divided in three layers, a shell of thickness 100 km, an ocean of about 200 km thick, and a solid core and the Love numbers formalism is used to compute the elastic deformation.  The internal structure of Titan is still under debate \\citep{Castillo10,Castillo12, McKinnon11, Fortes12}. The six different interior models selected here (see Table \\ref{table1}) are representative of the main categories of Titan's interior models. The librations computed for each model show that the libration amplitude for oceanic ($62-86$ m) or oceanless models ($51$ m) are of the same order at Titan's orbital frequency. We identify in the librations' Fourier decomposition two new frequencies resulting from the time-dependent inertia tensor in the external gravitational torque. The amplitudes of these librations are different for interior models with or without ocean. A reduction from $52$ to $30$ meters is obtained for the first term, and from $29$ to $18$ meters for the second one. These two long-period terms are interesting because they contain the information on the satellite internal structure.  Here, icy shell thicknesses of $100$, $75$ and $50$ kilometers have been investigated with a resulting libration amplitude at orbital frequency of the same order ($73.5$, $73.7$ and $61.6$ m, respectively). As shown by \\cite{VanHoolst13}, the sensitivity of librations is less dependent on the ice shell thickness when elasticity of the solid layers is taken into account.  The atmospheric torque CH12 obtained with a recent General Circulation Model from \\cite{LB12} presents a Saturnian semi-annual component ten times smaller than in the torque TO05 from \\cite{Tokano05}. At this frequency, the orbital perturbation induces a large librational response of about $560$ meters, and the atmospheric torque TO05 contributes for about $10\\%$ of the libration amplitude. The contribution is reduced to a few percent for the atmospheric torque CH12.  The librational response at Saturnian semi-annual frequency contain a small signature of the atmospheric circulation and the inner core inertia as shown in eq.(\\ref{atmosolution}). In addition, we found that for a Titan containing an internal ocean the elasticity strongly reduces the amplitude of high frequency librations and the resulting amplitude is at the same order as the amplitude for solid Titan. Finally we identified two new librations (with arguments $2\\varpi_8-2\\Omega_6-\\varpi_6$ and $2L_s-\\Omega_6-\\varpi_6$ ) at long period induced by the elasticity that contain a signature of the internal structure of a few ten meters. It is challenging to distinguish these small amplitudes librations with periods of about 30 and 50 years.   \\appendix"
    },
    "1402/1402.7049_arXiv.txt": {
        "abstract": "An expanding universe is not expected to have a static vacuum energy density. The so-called cosmological constant $\\Lambda$ should be an approximation, certainly a good one for a fraction of a Hubble time, but it is most likely a temporary description of a true dynamical vacuum energy variable that is evolving from the inflationary epoch to the present day. We can compare the evolving vacuum energy with a Casimir device where the parallel plates slowly move apart (``expand''). The total vacuum energy density cannot be measured, only the effect associated to the presence of the plates, and then also their increasing separation with time. In the universe there is a nonvanishing spacetime curvature $R$ as compared to Minkowskian spacetime that is changing with the expansion. The vacuum energy density must change accordingly, and we naturally expect $\\delta\\Lambda\\sim R\\sim H^2$. A class of dynamical vacuum models that trace such rate of change can be constructed. They are compatible with the current cosmological data, and conveniently extended can account for the complete cosmic evolution from the inflationary epoch till the present days. These models are very close to the $\\CC$CDM model for the late universe, but very different from it at the early times. Traces of the inherent vacuum dynamics could be detectable in our recent past. ",
        "introduction": "The cosmological term in Einstein's equations has traditionally been associated to a mysterious concept called the vacuum energy density: $\\rL = \\CC/(8\\pi G)$. With the advent of quantum theory and quantum field theory (QFT) this association has been even more troublesome and has led to the famous cosmological constant (CC) problem\\,\\cite{CCproblem}. The CC problem is the main source of headache for every theoretical cosmologist confronting his/her predictions with the observational value of $\\rL$\\,\\cite{CosmoObservations,CosmoObservations2,PLANCK1}. This is specially so after the discovery of a particle consistent with the Higgs boson at the LHC\\,\\cite{HiggsDiscovery}, as this may represent a serious endorsement to the notion of spontaneous symmetry breaking (SSB) and of electroweak vacuum energy in the context of the standard model (SM) of particle physics\\,\\cite{SM:EW,SM:QCD}. It is therefore more important than ever to properly address the notion of vacuum energy and its possible implications in cosmology. We will  elaborate on these important issues after considering the notion of dynamical vacuum energy and its interesting implications for a successful description of the entire cosmic history.  Despite its name, the possibility that the cosmological ``constant'' in Einstein equations is actually a time dependent quantity in cosmology, $\\CC=\\CC(t)$, is a most natural one, as it is difficult to conceive an expanding universe with a strictly constant vacuum energy density, $\\rL$. It would  be tantamount to accept that such a quantity has remained immutable since the origin of time without having any theoretical explanation for this fact. A smoothly evolving vacuum energy density $\\rL(t)=\\rL(H(t))$ that inherits its time-dependence from the Hubble rate $H=H(t)$ as the natural dynamical variable in the Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) background, is not only a qualitatively more plausible and intuitive idea, but is also suggested by fundamental physics, in particular by quantum field theory (QFT) in curved space-time\\,\\cite{BooksQFT} --  see also \\cite{JSP-CCReview13} for a recent review on CC aspects of it.   In order to implement the notion of time evolving vacuum energy density is not necessary to resort to \\textit{ad hoc} scalar fields, as usually done in the literature (e.g. in quintessence formulations and the like). A ``running'' term can be expected on very similar grounds as one expects (and observes) the running of couplings and masses with a physical energy scale in QFT. Furthermore, the experimental evidence that the equation of state (EOS) of the dark energy (DE) could be evolving with time/redshift (including the possibility that it might currently behave phantom-like) suggests that a time-variable term (possibly accompanied by a variable Newton's gravitational coupling too, $G = G(t))$ could account in a natural way for all these features. ",
        "conclusions": "We have discussed models with time evolving vacuum energy $\\rL(t)$ (including in some cases also variable gravitational coupling G) and assessed their impact in the cosmic history. We have emphasized that dynamical vacuum energy in an expanding universe is expected within QFT in curved spacetime\\,\\cite{JSP-CCReview13}. We have paid particular attention to models of running vacuum energy potentially connected with the possible quantum effects on the effective action of QFT in curved spacetime.  Specifically, we have shown that if the vacuum energy density evolves as a function of the form \\begin{equation}\\label{eq: GeneralRG} \\rL(t)=c_0+\\sum_{k=1} \\alpha_{k} H^{2k}(t)+\\sum_{k=1} \\beta_{k}\\dot{H}^{k}(t)\\,, \\end{equation} with $c_0\\neq 0$ (viz. an ``affine'' function constructed out of  powers of $H^2=\\left(\\dot{a}/a\\right)^2$ and $\\dot{H}=\\ddot{a}/a-H^2$, hence with an even number of time derivatives of the scale factor $a$), one can formulate a unified model of the cosmological evolution, compatible with the general covariance of the effective action, in which inflation is predicted, a correct transition (``graceful exit'') into a radiation phase can be naturally accommodated, and finally the late time cosmic evolution can also be successfully described.  \\vspace{0.25cm}  These kind of models, if we reparameterize them as a self-conserved DE fluid at fixed G (typically at $\\CC=0$, as done in quintessential philosophy), generally lead to a non-trivial effective equation of state.  In this picture, which is the common one in most studies of the dark energy, the running models can effectively appear as quintessence, and even as phantom energy or a mixture of both. This feature  can help explaining the observation of possible ``mirage transitions'' from quintessence to phantom, or vice versa, without invoking at all the existence of fundamental quintessence or phantom scalars fields. Our analysis serves also to illustrate the general fact that parametric degeneracies (or, as they are also called, cosmographic degeneracies) along with the unknown nature of dark energy makes it very tricky to work with parametric forms of the equation of state of dark energy fitting cosmological data. It turns out that the running vacuum models under consideration cannot be fitted to the most common forms of the equation of state of the dark energy. These are usually very simple, but prove unable to capture essential traits of more realistic models.  \\vspace{0.25cm} The recent observation of a Higgs-like boson particle at CERN's LHC\\,\\cite{HiggsDiscovery} reminded us that the vacuum energy (through the spontaneous symmetry breaking mechanism) can be a fully tangible concept in real phenomenology. Despite the usual difficulties with this concept, the Higgs discovery points once more to the notion of vacuum energy in quantum field theory. It should spur us to understand its meaning more deeply, rather than keep on blindly replacing it by a variety of other concepts, each one more exotic and fathomless than the previous.  In the meanwhile the effects associated to changes in the vacuum energy can be as real as they are in, say, the Casimir effect, which is sensitive only to tiny changes of the vibrational modes of the vacuum configuration. The possible analogy suggests that in the cosmological case the changes in the vacuum energy density are described in terms of mildly evolving functions of the expansion rate, specifically in the form of powers of $H^2$ and $\\dot{H}$, both being of order of the observed $\\CC$. These functions describe the time evolution of the spacetime curvature of our expanding universe and can ultimately be responsible for the currently observed cosmic acceleration. As in the Casimir effect, these terms are tracing the differences of the vibrational modes of the curved geometry of the expanding vacuum as compared to the Minkowski one, after we have renormalized away the bulk vacuum energy of the latter\\,\\cite{JSP-CCReview13}.  What is left is the measurable quantity that we may unnecessarily have called ``dark energy'' in the physical spacetime.  \\vspace{0.25cm}  As pointed out in our discussion, the ultimate origin of the ``dark energy'' might be connected to the covariant response of General Relativity impinged on the parameters of the gravitational sector, $\\CC$ and $G$, to the possible time evolution of the particle masses in the universe as triggered by the dynamics of particle physics \\,\\cite{FritzschSola2012,FritzschSola2014}. If so, this could be revealing an intimate connection between the quantum physics of the micro and the macro Cosmos.           \\begin{theacknowledgments} I have been supported in part by FPA2010-20807 (MICINN), Consolider grant CSD2007-00042 (CPAN) and by DIUE/CUR Generalitat de Catalunya under project 2009SGR502. I am grateful to the organizers of the workshop \\textit{Particle and Nuclear Physics at all scales, Astroparticle Physics and Cosmology}, held in Saint Petersburg, Russia, for the invitation and smooth organization of the conference, as well as for the partial support. It is a pleasure to thank S. Basilakos, H. Fritzsch, J.A.S. Lima and D. Polarski for the recent collaboration in some of the work presented here. Finally, I wish to express my gratitude to the Institute for Advanced Study at the Nanyang Technological University in Singapore for the hospitality while the writing of these proceedings was being finished. \\end{theacknowledgments}    \\newcommand{\\JHEP}[3]{ {JHEP} {#1} (#2)  {#3}} \\newcommand{\\NPB}[3]{{ Nucl. Phys. } {\\bf B#1} (#2)  {#3}} \\newcommand{\\NPPS}[3]{{ Nucl. Phys. Proc. Supp. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRD}[3]{{ Phys. Rev. } {\\bf D#1} (#2)   {#3}} \\newcommand{\\PLB}[3]{{ Phys. Lett. } {\\bf B#1} (#2)  {#3}} \\newcommand{\\EPJ}[3]{{ Eur. Phys. J } {\\bf C#1} (#2)  {#3}} \\newcommand{\\PR}[3]{{ Phys. Rep. } {\\bf #1} (#2)  {#3}} \\newcommand{\\RMP}[3]{{ Rev. Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\IJMP}[3]{{ Int. J. of Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRL}[3]{{ Phys. Rev. Lett. } {\\bf #1} (#2) {#3}} \\newcommand{\\ZFP}[3]{{ Zeitsch. f. Physik } {\\bf C#1} (#2)  {#3}} \\newcommand{\\MPLA}[3]{{ Mod. Phys. Lett. } {\\bf A#1} (#2) {#3}} \\newcommand{\\CQG}[3]{{ Class. Quant. Grav. } {\\bf #1} (#2) {#3}} \\newcommand{\\JCAP}[3]{{ JCAP} {\\bf#1} (#2)  {#3}} \\newcommand{\\APJ}[3]{{ Astrophys. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AMJ}[3]{{ Astronom. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\APP}[3]{{ Astropart. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AAP}[3]{{ Astron. Astrophys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\MNRAS}[3]{{ Mon. Not. Roy. Astron. Soc.} {\\bf #1} (#2)  {#3}} \\newcommand{\\JPA}[3]{{ J. Phys. A: Math. Theor.} {\\bf #1} (#2)  {#3}} \\newcommand{\\ProgS}[3]{{ Prog. Theor. Phys. Supp.} {\\bf #1} (#2)  {#3}} \\newcommand{\\APJS}[3]{{ Astrophys. J. Supl.} {\\bf #1} (#2)  {#3}}   \\newcommand{\\Prog}[3]{{ Prog. Theor. Phys.} {\\bf #1}  (#2) {#3}} \\newcommand{\\IJMPA}[3]{{ Int. J. of Mod. Phys. A} {\\bf #1}  {(#2)} {#3}} \\newcommand{\\IJMPD}[3]{{ Int. J. of Mod. Phys. D} {\\bf #1}  {(#2)} {#3}} \\newcommand{\\GRG}[3]{{ Gen. Rel. Grav.} {\\bf #1}  {(#2)} {#3}}"
    },
    "1402/1402.0175_arXiv.txt": {
        "abstract": "This work focuses on the study of the aqueous alteration process which acted in the main belt and produced hydrated minerals on the altered asteroids. Hydrated minerals have been found mainly on Mars surface, on main belt primitive asteroids and possibly also on few TNOs. These materials have been produced by hydration of pristine anhydrous silicates during the aqueous alteration process, that, to be active, needed the presence of liquid water under low temperature conditions (below 320 K) to chemically alter the minerals. The aqueous alteration is particularly important for unraveling the processes occurring during the earliest times of the Solar System history, as it can give information both on the asteroids thermal evolution and on the localization of water sources in the asteroid belt.  \\\\ To investigate this process, we present reflected light spectral observations in the visible region (0.4--0.94 $\\mu$m) of 80 asteroids belonging to the primitive classes C (prevalently), G, F, B and P, following the Tholen (1984) classification scheme. We find that about 65 \\% of the C-type  and all the G-type asteroids investigated reveal features suggesting the presence of hydrous materials, mainly a band centered around 0.7 $\\mu$m, while we do not find evidence of hydrated materials in the other low albedo asteroids (B, F, and P) investigated. \\\\ We combine the present observations with the visible spectra of asteroids available in the literature for a total of 600  primitive main belt asteroids.  We analyze all these spectra in a similar way to characterize the absorption band parameters (band center, depth and width) and spectral slope, and to look for possible correlations between the aqueous alteration process and the asteroids taxonomic classes, orbital elements, heliocentric distances, albedo and sizes. Our analysis shows that the aqueous alteration sequence starts from the  P-type objects, practically unaltered, and increases through the P $\\rightarrow$ F $\\rightarrow$ B $\\rightarrow$ C $\\rightarrow$ G asteroids, these last being widely aqueous altered, strengthening thus the results previously obtained by Vilas (1994). Around 50\\% of the observed C-type asteroids show absorption feature in the visible range due to hydrated silicates, implying that more than $\\sim$ 70\\% of them will have a 3 $\\mu$m absorption band and thus hydrated minerals on their surfaces, based on correlations between those two absorptions (Howell et al., 2011). \\\\ We find that the aqueous alteration process dominates in primitive asteroids located between 2.3 and 3.1 AU, that is at smaller heliocentric distances than previously suggested by Vilas et al. (1993). The percentage of hydrated asteroids is strongly correlated with their size. The aqueous alteration process is less effective for bodies smaller than 50 km, while it dominates in the 50--240 km sized primitive asteroids.  \\\\ No correlation is found between the aqueous alteration process and the asteroids albedo or orbital elements. Comparing the $\\sim$ 0.7 $\\mu$m band parameters of hydrated silicates and CM2 carbonaceous chondrites, the meteorites that have aqueous altered asteroids as parent bodies,  we see that the band center of meteorites is at longer wavelengths than that of asteroids. This difference on center positions may be attributed to different minerals abundances, and to the fact that CM2 available on Earth might not be representative of the whole aqueous altered asteroids population. ",
        "introduction": "The distribution of the asteroids taxonomic classes follows a radial structure varying with the  heliocentric distance: the inner and middle belt is dominated by differentiated bright asteroids (M, E, V, and prevalently S-types) that have experienced high temperatures in the past and that are constituted by volatile-poor silicates and metals, while dark primitive asteroids (B, P, D, and prevalently C-types), rich in carbonaceous chondrite-like materials and formed in a colder environment, dominate the outer belt to the Jupiter Trojans. This asteroid class distribution corresponds to a radial variation of the formation temperatures and to the presence of different mineralogical materials with increasing heliocentric distance. These materials have also been partially mixed in the past due to dynamical interaction of asteroids with each others and with planets, and altered over time by different processes such as collisions, aqueous alteration, and space weathering.  The aqueous alteration process acts on primitive asteroids (C, G, B, F and P-types, following the Tholen (1984) classification scheme) located mainly between 2.6 and 3.5 AU, in the so called zone of aqueous alteration (Vilas, 1994; Fornasier et al., 1999). The aqueous alteration process produces the low temperature ($<$ 320 K) chemical alteration of materials by liquid water which acts as a solvent and produces hydrated minerals such as phyllosilicates, sulfates, oxides, carbonates, and hydroxides. This means that liquid water was present in the primordial asteroids, produced by the melting of water ice by heating sources, very probably by  $^{26}$Al decay. \\\\ Reflectance spectroscopy of aqueous altered asteroids shows absorption features in the 0.6-0.9 and 2.5-3.5 micrometers regions, which are diagnostic of, or associated with, hydrated minerals.  The most prominent and unambiguous indicator of hydratation is the 3 $\\mu$m band (Lebofsky, 1980; Jones et al., 1990; Rivkin et al., 2002; Vilas, 1994, 2005; Howell et al., 2011; Takir \\& Emery, 2012),  associated with {\\it free} water molecules, and to OH ion bounded in the mineral crystal lattice; while in the visible range there are several bands, centered around 0.43 $\\mu$m, 0.60--0.65 $\\mu$m, 0.70 $\\mu$m and 0.80--0.90 $\\mu$m, attributed to charge transfer transitions in oxidized iron (Vilas et al., 1993, 1994; Vilas, 1994; Barucci et al., 1998; Fornasier et al., 1999).  The study of aqueous alteration is particularly important for unraveling the processes occurring during the earliest times in Solar System history, as it can give information both on the thermal processes and on the localization of water sources in the asteroid belt, and for the associated astrobiological implications.  Indeed it has been suggested (Morbidelli et al., 2000; Lunine, 2006) that the Earth current supply of water was delivered mostly by asteroids, not comets, some time following the collision that produced the Moon (which would have vaporized any of the water existing at that time). While asteroids and comets from the Jupiter-Saturn region were deemed to be the first water deliverers (when the Earth was half its size), the bulk of the water delivered was found to originate from a few planetary embryos, originally formed in the outer asteroid belt and accreted by the Earth at the final stage of its formation. Morbidelli et al. (2000) found that the measured amount of the D/H ratio in the water on Earth correlates very well with the ratio of D/H typical of water condensed in the outer asteroid belt. Alexander et al. (2012), analyzing the bulk hydrogen and nitrogen isotopic compositions of CI chondrites, suggest that these meteorites were the principal source of Earth  volatiles, while the D/H of all comets except Hartley 2 (Hartogh et al., 2012) are quite different from the Earth oceans value.  The snow line has been located within or nearby the asteroid belt (Lunine, 2006; Cyr et al., 1998), and, following the recent nebular kinetics models results, it may have migrated as the nebula evolved, sweeping across the entire asteroid belt (Dodson-Robinson et al., 2009). According to the latest dynamical models, it has also been suggested that part of the main belt primitive asteroids were formed either between the giant planets or in the Trans-neptunian region (Walsh et al., 2011), and then injected into the inner Solar System. \\\\ Water ice has been tentatively identified on the surfaces of the outer primitive asteroids Themis and Cybele (Rivkin and Emery, 2010; Campins et al., 2010; Licandro et al., 2011), and  some main belt asteroids, originated from the Themis family (Hsieh \\& Jewitt, 2006; Hsieh et al., 2009, 2012), were discovered to become active (they are also called main belt comets).  The fact that primitive asteroids had retained water ice in the past and may have enriched the Earth of water and possibly organic material, favoring the appearance of life on our planet, is of great interest for the scientific community. Indeed all of the major space agencies are planning sample return missions to primitive near Earth objects (NEO): NASA will launch OSIRIS-REx in 2016 to sample the asteroid (101955) Bennu, JAXA will launch Hayabusa-2 in 2014 to bring back material from (162173) 1999 JU3, and the mission MarcoPolo-R has been proposed to the European Space Agency (ESA) in the framework of the Cosmic Vision 2015--2025 program, for a sample return from (342843) 2008 EV5. \\\\ Some evidence of hydrated material has been found also in the outer Solar System, with the detection of peculiar absorption bands on the spectra of some TNOs that could be possibly  associated with the aqueous alteration process (Fornasier et al., 2004a, 2009; Lazzarin et al., 2003; de Bergh et al., 2004; Alvarez-Candal, 2008).  In this work we focus on spectroscopy in the visible region of low albedo asteroids. We present new visible spectra for 80 objects belonging to the C-complex and localized between 2.3 and 4.0 AU, with the aim to investigate the aqueous alteration process that has involved the  C, B, F, P, and G primitive classes. To do so, we have studied a larger sample including visible spectra from the literature for a total of 600 primitive main belt asteroids. We have analyzed all these data to spectrally characterize the bands  associated with the aqueous alteration process in the visible region (mainly the one centered around 0.7 $\\mu$m), and look for the relationships between this process and the heliocentric distance, albedo and diameter of the investigated objects. We have also compared the spectral parameters  associated with the 0.7 $\\mu$m band for asteroids and for the CM2 carbonaceous chondrites, which show evidence of aqueous altered materials on their surface  and in particular they show the 0.7 $\\mu$m band in their spectra. ",
        "conclusions": ""
    },
    "1402/1402.0519_arXiv.txt": {
        "abstract": "Measurements of absolute cross sections for the {\\it K}-shell photoionization of B-like atomic nitrogen ions were carried out utilizing the ion-photon merged-beam technique at the SOLEIL synchrotron radiation facility in Saint-Aubin, France.  High-resolution spectroscopy with E/$\\Delta$E $\\approx$ 13,500 was the maximum resolution achieved. We have investigated two photon energy regions: 404 eV -- 409 eV and 439 eV -- 442 eV.  Resonance peaks found in the experimental measured cross sections are compared with theoretical estimates from the multi-configuration Dirac - Fock, R-matrix and empirical methods, allowing identification of the strong 1s  $\\rightarrow$ 2p and the weaker 1s $\\rightarrow$ 3p resonances in the observed {\\it K}-shell spectra of this B-like nitrogen ion. ",
        "introduction": "X-ray spectra from {\\it XMM-Newton} may be utilized to characterize the interstellar medium (ISM), if accurate atomic {\\it K}-edge cross sections are available \\cite{McLaughlin2001, Brickhouse2010, Kallman2010, Quinet2011,Soleil2011,Soleil2013,McLaughlin2013}. Single and multiply ionization stages of C, N, O, Ne and Fe have been observed in the ionized outflow in  NGC 4051 measured with {\\it XMM-Newton} \\cite{Olge2004} in the soft x-ray region and low ionized stages of C, N and O have also been used in modelling x-ray emission from OB super-giants \\cite{Cassinelli1981}. Radiative/photo recombination of singly and doubly charged nitrogen ions also play an important role in the chemistry of the atmosphere of Titan \\cite{Dutuit2013}. Detailed photoionization models of the brightest knot of star formation in the blue compact dwarf galaxy Mrk 209 required abundances for ions of oxygen and nitrogen \\cite{Diaz2007}. The  XMM-{\\it Newton} X-ray spectra of WR 1 is rich in nitrogen ions  \\cite{Ignace2003} and photoionization cross section data and abundances for carbon, nitrogen, and oxygen in their various stages  of ionization are essential for photoionization models applied to the plasma modelling in a variety of planetary nebulae \\cite{Bohigas2008}. In the present study we focus our  attention on obtaining detail spectra on the doubly ionized nitrogen ion N$^{2+}$ (N III) in the vicinity of its {\\it K} - edge.  Photoionization (PI) cross sections used for  the modelling of astrophysical phenomena have mainly been provided by theoretical methods, due to limited experimental data being available. Major effort has gone into improving the quality of calculated data using state-of-the-art theoretical methods. Recent advances in the determination of atomic parameters for modeling {\\it K} lines in cosmically abundant elements have been reviewed by Quinet and co-workers \\cite{Quinet2011}.  Absolute experimental {\\it K}-shell photoionization cross section results have been obtained by various groups on a variety of atoms and ions of astrophysical interest; He-like Li$^{+}$ \\cite{Scully2006,Scully2007,DPI2013}, Li-like  B$^{2+}$ \\cite{Mueller2010}, C$^{3+}$  \\cite{Mueller2009}, N$^{4+}$ \\cite{Soleil2013}, Be-like B$^{+}$ \\cite{Mueller2014}, C$^{2+}$ \\cite{Scully2005}, N$^{3+}$ \\cite{Soleil2013}, B-like C$^{+}$ \\cite{Schlachter2004}, C-like N$^{+}$ \\cite{Soleil2011}, N-like O$^{+}$ \\cite{Kawatsura2002}, F-like Ne$^{+}$ \\cite{Yamaoka2001}, neutral nitrogen \\cite{McLaughlin2011} and oxygen \\cite{Krause1994,Menzel1996,Stolte1997,Stolte2013}.  Recent studies on {\\it K}-shell photoionization cross sections calculations for neutral nitrogen and oxygen showed excellent accord with high resolution measurements made at the Advanced Light Source (ALS) radiation facility \\cite{McLaughlin2011,Stolte2013} as have similar cross section calculations on singly and multiply ionized stages of atomic nitrogen compared with high resolution measurements at the SOLEIL synchrotron facility \\cite{Soleil2011,Soleil2013}. The majority of the high-resolution experimental data from third generation light sources have been shown to be in excellent agreement with the state-of-the-art R-matrix method \\cite{rmat,codes}  and with other modern theoretical approaches.  The present  investigation for this proto-type B-like atomic nitrogen ion gives accurate values of photoionization cross sections produced by x-rays in the vicinity of the {\\it K}-edge, where strong n=2 inner-shell resonance states of N$^{2+}$ are observed. This work compliments our previous studies on {\\it K}-shell photoionization of singly and multiply ionized atomic nitrogen ions \\cite{Soleil2011,Soleil2013} in the vicinity of the K-edge. Previous experimental studies on B-like atomic nitrogen have been performed only in the valence shell region \\cite{bizau2004,bizau2005}. To date no experimental studies for the {\\it K}-edge region have been reported  in the literature. For the  B-like ions N$^{2+}$, O$^{3+}$ and F$^{4+}$, absolute experimental cross section measurements in the near threshold region were made by  Bizau and co-workers \\cite{bizau2004,bizau2005}. Close-coupling calculations performed by Li-Guo and Xin-Xiao \\cite{china2009} based on the R-matrix formalism \\cite{Burke2011} obtained excellent agreement with the experimental work of Bizau and co-workers \\cite{bizau2005}. In this near threshold region it was necessary to include both the ground state and metastable excited states in the theoretical work in order to achieve suitable agreement with experiment.  We follow a similar approach here for the {\\it K}-shell energy region. N$^{2+}$ ions produced in the SOLEIL synchrotron radiation experiments are not purely in their ground state. {\\it K}-shell photoionization contributes to the ionization balance in a more complicated way than outer shell photoionization. In fact {\\it K}-shell photoionization when followed by Auger decay couples three or more ionization stages instead of two in the usual equations of ionization equilibrium \\cite{Petrini1997}.  The 1s $\\rightarrow$ 2p photo-excitation process on the $\\rm 1s^22s^22p~^2P^o$ ground-state of B-like nitrogen ion is, $$ h\\nu + {\\rm N^{2+}(1s^22s^22p~^2P^o)}  \\rightarrow  {\\rm N^{2+} ~ (1s2s^2 \\,2p^2[^3P, ^1D, ^1S] ~^2S, ^2P, ^2D) } $$ which can decay to $$ \\swarrow \\quad \\searrow $$ $$ {\\rm  N^{3+}~ (1s^22s^2~^1S) + e^- ({\\it k^2_{\\ell}}),} \\; \\; {\\rm or} \\; \\; {\\rm  N^{3+}~ (1s^22s~ np~^{1}P^{\\circ}) + e^- ({\\it k^2_{\\ell}}),} $$ where $k^2_{\\ell}$ is the outgoing energy of the continuum electron with angular momentum $\\ell$. Experimental studies on this doubly ionized atomic nitrogen ion, in its ground state $\\rm 1s^22s^22p~^2P^o$, are also hampered by metastable states present in the parent ion beam. In the present experimental studies performed at the SOLEIL radiation facility, N$^{2+}$ ions are produced in the gas-phase using an Electron-Cyclotron-Resonance-Ion-Source (ECRIS) so the metastable state $\\rm 1s^22s2p^2~^4P$  can be present in the parent ion beam.  For the $\\rm 1s^22s2p^2~^4P$ metastable state, auto-ionization processes occurring by the 1s $\\rightarrow$ 2p photo-excitation process  are mainly; $$ h\\nu + {\\rm N^{2+}(1s^22s2p^2~^4P)} $$ $$ \\downarrow $$ $$ {\\rm  N^{2+} [1s2s[^{1,3}S]\\,2p^3(^4S^o,^2D^o,^2P^o)]^4S^o,^4P^o,^4D^o} $$ $$ \\downarrow $$ $$ {\\rm N^{3+} (1s^22s~np~ ^3P^o) + e^- ({\\it k^2_{\\ell}}).} $$ In the vicinity of the {\\it K}-edge our current investigations appear to be the first time this B-like system has been studied experimentally for such an energy region.  State-of-the-art  {\\it ab initio} calculations for Auger inner-shell processes were first carried out on this B-like system by Petrini and de Ara\\'ujo \\cite{Petrini1997} using the R-matrix method \\cite{rmat} and followed a similar procedure to the work on {\\it K}-shell studies for the Be-like B$^+$ ion \\cite{Petrini1981}. Stoica and co-workers \\cite{Petrini1998} noted that once  the 1s-hole was created in the ions, by single photoionization, with simultaneous shake-up and shake-off processes, Auger decay populates directly excited states of the residual ions, which then produces UV lines. This work was further extended by  Garcia and co-workers \\cite{Witthoeft2009}, using the optical potential method within the Breit-Pauli R-matrix formalism \\cite{rmat,codes,damp,Burke2011}, for photoionization of the ground state only, along the nitrogen iso-nuclear sequence. Garcia and co-workers  \\cite{Witthoeft2009}  pointed out that the earlier central field calculations \\cite{Verner1993,Verner1995,rm1979} did not account for the strong autoionizing resonance features that dominate the cross sections near the {\\it K}-edge.  In the present study we compare our  theoretical cross section results from the multi-configuration Dirac Fock (MCDF) and R-matrix methods with previous theoretical results \\cite{Chen1987,Chen1988,Witthoeft2009} and with the current experimental measurements made at the SOLEIL synchrotron radiation facility. Detailed measurements of the  {\\it K}-shell  single photon ionization  cross sections for B-like nitrogen ions, were made in the 404 eV -- 409 eV region (where strong peaks were observed) and in the photon energy range 439 eV -- 442 eV. The results for resonance energies and Auger widths are compared with detailed theoretical predictions made using the MCDF \\cite{Bruneau1984}, R-matrix \\cite{Burke2011} and the SCUNC empirical methods \\cite{Sakho2013a,Sakho2013b}. The theoretical predictions assist in the identification and characterization of the strong $\\rm 1s \\rightarrow 2p$ and  the weaker $\\rm 1s \\rightarrow 3p$ resonances observed in the B-like nitrogen spectra. The current study gives absolute values (experimental and theoretical) for PI cross sections along with n=2 and n=3 inner-shell resonance energies, natural line widths and resonance strengths, for the situation of a photon interacting with the ground $\\rm 1s^22s^22p~^2P^o$,  and metastable $\\rm 1s^22s2p^2~^4P$ states of the N$^{2+}$ ion.  In section 2 we briefly outline the experimental procedure used and section 3 presents the theoretical procedures used. Section 4 gives a discussion of our experimental and theoretical results. Finally in section 5 conclusions are drawn from the present study. ",
        "conclusions": "\\label{sec:Conclusions} {\\it K}-shell photoionization of B-like nitrogen ions, N$^{2+}$, has been investigated using state-of-the-art experimental and theoretical methods. High-resolution spectroscopy was able to be achieved with E/$\\Delta$E = 13,500 at the SOLEIL synchrotron radiation facility, in Saint-Aubin, France, covering the photon energy ranges; 404 eV -- 409 eV and 439 eV -- 442 eV. Several strong peaks are found in the cross sections in the energy regions,  404 eV - 409 eV and  439 eV -- 442 eV which are identified as the 1s $\\rightarrow$ 2p and  1s $\\rightarrow$ 3p transitions in the N$^{2+}$ {\\it  K}-shell spectrum and assigned spectroscopically with their resonance parameters tabulated in Tables \\ref{reson} and \\ref{reson2}. For the observed peaks, suitable agreement is found between the present theoretical and experimental results both on the photon-energy scale and on the absolute cross-section scale for this prototype B-like system. The strength of the present study is the high resolution of  the spectra along with theoretical predictions made using state-of-the-art MCDF, R-matrix with pseudo-states and empirical methods. The present results have been compared with high resolution  experimental measurements made at the SOLEIL synchrotron radiation facility and with other theoretical methods so would be suitable to be incorporated into astrophysical modelling codes like CLOUDY \\cite{Ferland1998,Ferland2003}, XSTAR \\cite{Kallman2001} and AtomDB \\cite{Foster2012} used to numerically simulate the thermal and ionization structure of ionized astrophysical nebulae.  \\ack The experimental measurements were performed on the PLEIADES beamline \\cite{Pleiades2010,Miron2013}, at the SOLEIL Synchrotron radiation facility in Saint-Aubin, France.  The authors would like to thank the SOLEIL staff and, in particular those of the PLEIADES beam line for their helpful assistance, C Miron,  Xiao-Jing Liu, E Roberts and C Nicolas. B M McLaughlin acknowledges support from the US National Science Foundation through a grant to ITAMP at the Harvard-Smithsonian Center for Astrophysics, the RTRA network {\\it Triangle de le Physique} and a visiting research fellowship from Queen's University Belfast. I Sakho acknowledges the hospitality of the Universit\\'{e} Paris-Sud and support from the University Assane Seck of Ziguinchor during a recent visit. We thank Dr John C Raymond and Dr Randall K Smith at the Harvard-Smithsonian Center for Astrophysics for discussions on the astrophysical applications and Dr Javier Garcia for numerical values of the cross sections  \\cite{Witthoeft2009}. The computational work was carried out at the National Energy Research Scientific Computing Center in Oakland, CA, USA, the  Kraken XT5 facility at the National Institute for Computational Science (NICS) in Knoxville, TN, USA and at the High Performance Computing Center Stuttgart (HLRS) of the University of Stuttgart, Stuttgart, Germany. We thank Stefan Andersson from Cray Research for his advice and assistance with the implementation and optimization of the parallel R-matrix codes on the Cray-XE6 at HLRS. The Kraken XT5 facility is a resource of the Extreme Science and Engineering Discovery Environment (XSEDE), which is supported by National Science Foundation grant number OCI-1053575."
    },
    "1402/1402.5420_arXiv.txt": {
        "abstract": "We classify the spectral energy distributions (SEDs) of 431,038 sources in the 9 sq. deg \\Bootes\\ field of the NOAO Deep Wide-Field Survey (NDWFS).  There are up to 17 bands of data available per source, including ultraviolet (\\emph{GALEX}), optical (NDWFS), near-IR (NEWFIRM), and mid-infrared (IRAC/MIPS) data, as well as spectroscopic redshifts for $\\sim$20,000 objects, primarily from the AGN and Galaxy Evolution Survey (AGES).  We fit galaxy, AGN, stellar, and brown dwarf templates to the observed SEDs, which yield spectral classes for the Galactic sources and photometric redshifts and galaxy/AGN luminosities for the extragalactic sources.  The photometric redshift precision of the galaxy and AGN samples are  $\\sigma/(1+z)=0.040$ and $\\sigma/(1+z)=0.169$, respectively, with the worst 5\\% outliers excluded.  Based on the \\chisqred\\ of the SED fit for each SED model, we are able to distinguish between Galactic and extragalactic sources for sources brighter than $I=23.5$.   We compare the SED fits for a galaxy-only model and a galaxy+AGN model.  Using known X-ray and spectroscopic AGN samples, we confirm that SED fitting can be successfully used as a method to identify large populations of AGN, including spatially resolved AGN with significant contributions from the host galaxy and objects with the emission line ratios of ``composite'' spectra.  We also use our results to compare to the  X-ray, mid-IR, optical color and emission line ratio selection techniques.   For an $F$-ratio threshold of $F>10$ we find 16,266 AGN candidates brighter than $I=23.5$ and a surface density of $\\sim1900$ AGN deg$^{-2}$. ",
        "introduction": "\\label{sec:Intro} Active Galactic Nuclei (AGN) present a wide variety of observational properties, which can be exploited to identify them in large surveys using a range of selection techniques.  However, most selection methods are based on the characteristics in a limited wavelength region.  This leads to selection methods that are  incomplete for some AGN luminosities, classes, or redshifts.  For example, the early single (optical) color searches have evolved into multi-color optical surveys such as the Sloan Digital Sky Survey (SDSS), which has successfully identified $\\sim10^{5}$ quasars based largely on their optical colors \\citep{Richards2002, Richards2005, Schneider2007, Schneider2010}.  The color selection methods of \\citet{Richards2002} achieved a completeness of $\\sim$90\\% for unobscured, unresolved AGN with $i<19$, $z<5$, and emission dominated by the AGN, with an efficiency (number of quasars/number of candidates) of $\\sim$65\\%.  However, using $ugriz$ colors also presents difficulties in distinguishing between F stars and quasars at $z\\sim2$.  Since quasar activity is known to peak near this redshift \\citep{Richards2006, Assef2011}, this presents a problem for studying the accretion history of the universe.   Using broadband optical colors to find AGN also relies on the roughly universal UV/optical power-law shape of direct emission from an accretion disk to distinguish AGN from the primarily photospheric emission of stellar populations in this wavelength range.  Obscured, type-2 AGN are more difficult to identify using optical colors, due both to the distortion of the shape of the AGN continuum and the increasingly significant host galaxy contribution.  Similarly, low accretion rate AGN are dominated by their host emission \\citep[e.g.][]{Hopkins2009}.   X-ray emission is sensitive to both unobscured and moderately obscured AGN, with the added advantage that X-ray emission almost always implies the presence of AGN except at low luminosities ($L_{X}\\lesssim10^{40}$ erg s$^{-1}$) where binary contributions become important \\citep{Shapley2001}.  However, \\emph{Chandra} and \\emph{XMM-Newton} surveys are biased against highly obscured AGN whose soft X-rays are absorbed by large columns of gas and dust \\citep[e.g.,][]{Polletta2006}.  Hard X-ray observations largely avoid this problem, but until the recent launch of \\emph{NuSTAR} \\citep{Harrison2013,Alexander2013}, with its greater sensitivity to hard X-rays, there were only shallow surveys for the highly obscured, Compton-thick AGN, such as the \\emph{Swift} BAT Survey \\citep{Tueller2008}.  Radio surveys can also be used to search for AGN; however, typical radio surveys are limited to a very small population of AGN, since only $\\sim$10\\%-20\\% are radio-loud \\cite[e.g.][]{Kellermann1989, Urry1995, Stern2000, Ivezic2002}.   Mid-infrared (mid-IR) surveys are sensitive to both type-1 and type-2 AGN.  The accretion disk power law spectrum extends into the rest-frame near-IR and is less affected by extinction.  Additionally, at rest-frame mid-IR and longer wavelengths, there is emission by hot dust in the ``torus'' where the UV and optical emission from the accretion disk is absorbed and reprocessed \\citep[e.g.,][]{Sanders1989,Netzer2007}.  These contributions roughly produce a power-law spectrum in the mid-IR that is  easily distinguished from the infrared spectra of stars or relatively normal galaxies which are essentially falling Rayleigh-Jeans spectra regardless of stellar population or age at these wavelengths.  At lower redshifts ($z\\lesssim1.5$) mid-IR selection triggers on dust emission, while at higher redshifts ($z\\gtrsim1.5$) it really triggers on the direct emission from the disk.  As a result, many AGN occupy a region of mid-IR color space that is well separated from that of galaxies and stars \\cite[e.g.,][]{Lacy2004, Stern2005, Alonso2006, Donley2007}, which was initially demonstrated using data from the four bands (3.6\\micron\\ to 8.0\\micron) of \\emph{Spitzer}/IRAC.  The \\emph{Wide-Field Infrared Survey Explorer} \\cite[\\emph{WISE};][]{Wright2010} can be used to extend these methods over the full sky \\citep{Mateos2012, Stern2012, Assef2013, Yan2013}.  However,  mid-IR color selection methods fail when the host galaxy contribution becomes large, or a strong H$\\alpha$ line lies in the 3.6\\micron\\ band at $z\\sim4$ and distorts the colors \\citep{Assef2010}.  For quasars that are as faint as W2$\\sim$15.5 in the WISE 4.6\\micron\\ band, galaxy contamination becomes significant and mid-IR color selection becomes more difficult \\citep{Assef2013}.  Each technique for identifying AGN has its own set of selection effects and it is important to understand the differences between AGN samples selected with the various methods.  For example, AGN that are detected using X-ray, radio, or mid-IR colors have different characteristic Eddington ratios, clustering properties, and host galaxy morphologies and host galaxy masses. \\citet{Hickox2009}  studied a sample of $\\sim$600 AGN at $0.25<z<0.8$ with spectroscopic redshifts from the AGES and Galaxy Evolution Survey \\cite[AGES;][]{Kochanek2012}, and radio, X-ray, and mid-IR data from the Westerbork Synthesis Radio Telescope, the \\emph{Chandra} X\\Bootes\\ Survey, and the \\emph{Spitzer} IRAC Shallow Survey, respectively.  With modest overlap between the radio, X-ray, and mid-IR AGN samples, \\citet{Hickox2009} found that radio AGN tend to be optically red and have massive host galaxies (i.e. red sequence galaxies), while X-ray AGN span a large range of host masses and colors but show a peak in the ``green valley'', and mid-IR AGN tend to be slightly optically bluer than the X-ray AGN.  \\citet{Griffith2010} arrived at a similar conclusion by examining the host galaxy morphologies of radio, mid-IR, and X-ray selected AGN, and found that radio-selected AGN tend to have early-type hosts, while mid-IR and X-ray selected AGN are more likely to have disky host galaxies.  This may represent a sequence in evolution of the quasar phase, which is thought to be triggered by a burst of star formation that funnels gas into the central engine, which ultimately quenches the star formation, transforming the optical colors from blue to red \\citep{Hopkins2007, Cardamone2010, Povic2012}.   Traditionally, extragalactic surveys have either been wide and shallow, or narrow and deep.  Wide/shallow surveys such as the 2 Micron All-Sky Survey \\cite[2MASS;][]{Skrutskie2006}, the Sloan Digital Sky Survey \\cite[SDSS;][]{Abazajian2009}, and the 2 Degree Field Galaxy Redshift Survey \\cite[2dF;][]{Colless2001} have allowed studies of large numbers of nearby galaxies and contributed significant discoveries about galaxy properties in the low redshift universe and luminous AGN at all redshifts \\citep[e.g.][]{Richards2005,Huchra2012}.  Examples of deep/narrow surveys include the Great Observatories Origins Deep Survey \\cite[GOODS;][]{Dickinson2003} and the Cosmological Evolution Survey \\cite[COSMOS;][]{Scoville2007}.  GOODS covers $\\sim$300 arcmin$^{2}$ with deep observations from \\emph{Spitzer}, \\emph{Hubble}, \\emph{Chandra}, \\emph{Herschel} and \\emph{XMM-Newton}, as well as ground-based facilities.  COSMOS is a wider survey, covering $\\sim$2  deg$^{2}$ with imaging from both space-based and ground-based telescopes, including \\emph{Hubble}, \\emph{Spitzer}, \\emph{GALEX}, \\emph{Chandra}, \\emph{NuSTAR}, Subaru, VLT, and others.   The relative rarity of AGN makes it difficult however to fully characterize AGN in deep fields or to fully account for faint AGN in shallow fields.  While these (and other) extragalactic surveys have provided a tremendous contribution to our knowledge of both nearby and distant galaxy populations, there have been few surveys on $\\sim$10 deg$^{2}$  scales.  The NOAO Deep Wide-Field Survey \\citep[NDWFS;][]{Jannuzi1999} in the 9 deg$^{2}$ \\Bootes\\ field allows us to explore a cosmologically significant volume with data that is both deep and extensive in wavelength coverage.  The large survey volume means that studies in the \\Bootes\\ field are less affected by cosmic variance or small number statistics for rare, faint objects that would not be found in either deep/narrow or wide/shallow surveys.  In addition to near-IR and optical data from NDWFS, there is a wealth of ancillary data in the \\Bootes\\ field.  This includes additional optical ground-based data \\citep{Cool2007,Bian2013}, X-ray data from the \\emph{Chandra X-ray Observatory}  \\citep{Murray2005, Kenter2005, Brand2006}, UV data from the \\emph{Galaxy Evolution Explorer} \\cite[\\emph{GALEX};][]{Hoopes2004}, and mid-IR data from the \\emph{Spitzer Space Telescope} \\citep{Eisenhardt2004, Ashby2009, Jannuzi2010}.   There is also spectroscopic data from the AGN and Galaxy Evolution Survey \\cite[AGES;][]{Kochanek2012}, which provides redshifts for $\\sim$18,000 galaxies ($I<20$) and $\\sim$4,700 AGN candidates ($I<22.5$).   In this paper, we examine and classify the UV to mid-IR spectral energy distributions (SEDs) of 431,038 sources in the \\Bootes\\ survey region.  By fitting galaxy, AGN, and stellar templates to the observed SEDs, we are able to isolate AGN candidates from galaxy and stellar populations using a much broader wavelength range than the common AGN selection methods discussed previously.  In \\S\\ref{sec:data} we outline the data that are used in this work.  In \\S\\ref{sec:analysis-SEDfitting} and \\S\\ref{sec:analysis-droppedfilters} we compare photometric to spectroscopic redshifts and discuss in detail the SED fitting procedure and how our final sample of SED fits were assembled.  In \\S\\ref{sec:analysis-galorexgal} we illustrate the separation of Galactic and extragalactic sources based on \\chisqred\\ statistics and we investigate the difficulty of classifying sources at faint magnitudes.   Section \\S\\ref{sec:AGNcandidates} presents how AGN candidates may be selected based on the results of their SED fits as galaxy and galaxy+AGN models and the sensitivity of this SED selection to the relative strength of the underlying host galaxy. In \\S\\ref{sec:compare}, we compare the SED selection of AGN candidates to mid-IR and optical color selections. In \\S\\ref{sec:surface_density} we compare the surface densities of different AGN samples. Finally in \\S\\ref{sec:future} we summarize our most important results and comment on future work.   Throughout this work we assume a flat $\\Lambda$CDM cosmology with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.28$, and $\\Omega_{\\Lambda}=0.72$. ",
        "conclusions": "\\label{sec:future}  Using photometric data that ranges from the far-UV to the mid-IR, we fit galaxy, AGN, stellar, and brown dwarf SED models to 431,038 sources in the \\Bootes\\ NDWFS field.  The photometric separation of stars and galaxies compares well with morphological separation and star and galaxy surface densities at all magnitudes.  Comparing to the over 20,000 available spectroscopic redshifts, we find photometric redshift dispersions of $\\sigma/(1+z)$=0.040 and $\\sigma/(1+z)=0.169$ for the galaxy and AGN samples, respectively, after clipping the worst 5\\% of sources.  In practice, the AGES survey \\citep{Kochanek2012} obtained spectroscopic redshifts for a large fraction of the most problematic sources, the luminous broad line AGN, so the photometric redshift dispersion of the sources without spectra should be lower.  We estimated the likelihood of an AGN component based on F-test comparisons of the fits.  The results for all the sources are reported in Table~\\ref{table:results}.  When we examine the distribution of F-ratios as a function of morphology and host luminosity fraction, we see the expected trends that high F-ratio sources tend to be more point-like and have smaller host galaxy contributions to their SEDs.  X-ray sources show a broader distribution of F-ratios and host galaxy contributions. We also examined the distribution of sources relative to the \\citet{Stern2005}, \\citet{Lacy2004} and \\citet{Donley2012} mid-IR selection criteria.  As a comparison, 62\\%, 70\\% and 35\\% of the X\\Bootes\\ X-ray sources satisfy these criteria.  There is a clear trend of the sources moving into these selection regions as $F$ increases, with 57\\%, 75\\% and 32\\% of the sources with F-ratios $F>10$ for adding an AGN component falling in these mid-IR selection regions.  Sources with $F>40$ lie almost exclusively in these regions.  The $F \\sim 10$ sources have significant host contributions, and many lie just bluewards of these selection regions in their $[3.6]-[4.5]$ colors. They partially overlap the mid-IR color distribution of X-ray sources, but they extend towards redder $[5.8]-[8.0]$ colors.  Their mid-IR color distribution is very similar to that of narrow line sources with ``composite'' line ratios indicative of a mixture of star formation and AGN activity, which is also consistent with their redder $[5.8]-[8.0]$ colors.  Similarly, we can use the F-ratio distribution to examine how AGN extend outwards from the SDSS optical color selection regions as the host galaxy contribution becomes more important.  These results strongly suggest that a sample of $F>10$ sources can provide a more complete inventory of AGN activity than any of the methods restricted to limited wavelength regimes.  Like any method, it is not perfect. There are clearly a minority of X-ray AGN with such low F-ratios that they cannot be identified based on their UV-IR SEDs with an acceptable false positive rate. What needs to be calibrated at this point are the false positive rates as a function of the apparent significance of the AGN contribution.  The surface densities of $F>10$ sources as compared to all extragalactic sources and other AGN samples suggests that the false positive rate should be acceptably low, but this needs to be tested with spectroscopy.  For $F\\sim 5$, the false positive rate is likely unacceptably high.  The problem for spectroscopy is that many redshift ranges will lack the emission line diagnostics needed to classify the nature of the source.  Most surveys of evolution separate the study of galaxies and AGN because they have difficulties tracking the populations which are strong mixtures of both.  This large scale decomposition of the SEDs into host and AGN components provides a means of studying the co-evolution of these populations with relatively reliable estimates of the two components separately.  Even when the models have only photometric redshifts, \\citet{Assef2010} found that the estimates of the host galaxy luminosity fraction were relatively robust.  In particular, it should be feasible to examine the duty cycle of AGN activity in galaxies as a function of redshift.  For example, if the false positive rates for AGN activity can be calibrated as a function of $F$ using modest spectroscopic samples, then the full survey sample can be used in studies of the evolution of galaxies and AGN.  These very broad baseline SED models should be comparably ``stable'' to surveys using large numbers of narrower filters over smaller wavelength ranges \\citep[e.g.,][]{Geach2008,Abramo2012}.  Instead of trying to better identify spectral breaks or strong emission lines, the broad structure of galaxy SEDs with a peak in the near-IR allows robust photometric redshifts and the structure of the UV and mid-IR tails of the distribution provides a robust probe of star formation and AGN activity.  This approach should be particularly valuable for integrating wide area, multi-wavelength surveys in the ultraviolet (\\emph{GALEX}), optical (SDSS, DES\\footnote{http://www.darkenergysurvey.org/}, Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu/public/}, LSST\\footnote{http://www.lsst.org/}), near-IR (VISTA, \\citet{Emerson2006}, \\emph{EUCLID}\\footnote{http://sci.esa.int/euclid/45403-mission-status/}, \\emph{WFIRST}\\footnote{http://wfirst.gsfc.nasa.gov/}) and mid-IR (\\emph{WISE}, \\citet{Wright2010}, \\emph{Akari} \\citep{Murakami2007})."
    },
    "1402/1402.5328_arXiv.txt": {
        "abstract": "{ We propose a class of models of supersymmetric hilltop inflation (also called ``new inflation'') where the initial conditions of the inflaton close to the hilltop are generated through ``matter field preinflation''. This is achieved via a coupling term between the inflaton and matter fields (i.e. Standard Model fields or a right-handed neutrino). The same coupling also opens up a decay channel for the inflaton into Standard Model fields, which allows efficient reheating of the universe. We discuss the multifield dynamics of the inflaton and matter fields during inflation using the $\\delta N$ formalism and show under which conditions the model effectively reduces to single-field hilltop inflation during the last 60 e-folds. We also study perturbative reheating through the matter-inflaton coupling for a specific example where the matter field is identified with a right-handed (s)neutrino, and demonstrate that in this case the model can generate the observed baryon asymmetry through nonthermal leptogenesis. }       \\begin{document} ",
        "introduction": "Inflation provides a successful paradigm for the early universe, capable not only of resolving the flatness and horizon problems of big bang cosmology, but also of generating primordial perturbations which eventually develop into the present day inhomogeneities of the universe and can be probed by observations of the CMB and large scale structures.  The Planck satellite has recently observed the temperature anisotropies of the CMB with unprecedented precision~\\cite{Ade:2013zuv,Ade:2013uln,Ade:2013ydc}. The new data has brought further confirmation of inflation and provided useful information towards discriminating between different models of inflation. In particular, the non-Gaussianity parameter has been constrained to $f_{\\text{NL}} = 2.7 \\pm 5.8$ at $68\\%$ C.L., which is consistent with single-field slow-roll inflation. The spectral index $\\ns$ has been measured to be $\\ns = 0.9603\\pm0.0073$ at $68\\%$ C.L., which constrains the possible shapes of the inflaton potential, and the upper bound on the tensor-to-scalar ratio has been tightened to $r < 0.11$ at $95\\%$ C.L., which already excludes some large-field models of inflation.  Among the various inflation scenarios, hilltop inflation \\cite{hilltop,newinf,newinfSG,Senoguz:2004ky} (also called ``new inflation'') can lead to a primordial spectrum close to Planck's best-fit value \\cite{Ade:2013uln}. In hilltop inflation, the inflaton field rolls down from a maximum of a very flat potential which is shaped like a plateau. Such potentials can arise when a phase transition takes place, for example associated with the breaking of a GUT symmetry \\cite{newinf} or a flavour symmetry \\cite{Antusch:2008gw}, which allows for possible close connections to particle physics.  However, given that the inflaton field $\\phi$ has to start its dynamics close to the maximum of its potential, the obvious question arises of how the field initially gets there. Many of the scenarios proposed to resolve this initial condition problem introduce a preceding period of slow-roll inflation (called ``preinflation'') driven by a preinflaton $\\chi$. During this phase of preinflation, the inflaton field $\\phi$ is forced to the maximum of its potential, which dynamically generates the initial conditions required for hilltop inflation in $\\phi$. Usually, this involves adding a new sector of singlet fields which are decoupled from the visible sector (e.g. in \\cite{Izawa:1997df,Senoguz:2004ky,smoothPreinflation}). A similar function is performed by a coupling of the inflaton to the energy density as proposed in \\cite{symmetron}. Other approaches to the initial conditions problem include \\cite{otherICs}.  In this paper, we show how preinflation can be realized in supersymmetric models using only a coupling of the inflaton to ``matter fields'', which we define as Standard Model fields or right-handed neutrinos (and their superpartners). In the stage of preinflation, the scalar component of a matter superfield (or alternatively of a D-flat combination of charged matter superfields) takes the role of the preinflaton and generates a mass term for the inflaton which leads to the right initial conditions for hilltop inflation. After hilltop inflation, the same coupling to matter fields allows the inflaton to decay and reheat the universe.  We also study the effects of the preinflaton matter field on the primordial spectrum, and we illustrate how reheating can occur via the inflaton-matter coupling for the example of a right-handed sneutrino preinflaton.  The paper is organized as follows: in sections~\\ref{sec:idea} and~\\ref{sec:model} we discuss qualitatively the main ideas of the paper and present the model. We then calculate in section~\\ref{sec:ics} the likely initial conditions that arise from quantum fluctuations near the instability point. Using these initial conditions, in section~\\ref{sec:numericalResults} we numerically calculate the spectrum of perturbations from inflation using the $\\delta N$ formalism (briefly introduced in section~\\ref{sec:deltaN}), confirming that single-field hilltop inflation preceded by matter field preinflation is a generic outcome of our model setup for a large region of parameter space. Finally, in section~\\ref{sec:reheating} we demonstrate for a specific example model that the inflaton-matter coupling can provide a viable channel for perturbative reheating and leptogenesis. ",
        "conclusions": "In this paper, we have proposed a new class of models of hilltop inflation where the initial conditions of the inflaton close to the hilltop are generated from ``matter field preinflation''. This is achieved via a coupling term between the inflaton and a matter field (cf. second term in eq.~\\eqref{wgen}), in addition to the usual term for supersymmetric hilltop inflation. The same coupling term also opens up a decay channel for the inflaton into Standard Model fields, which allows efficient reheating of the universe.  Our mechanism works as follows: During preinflation, the scalar component of the matter superfield (or of a D-flat combination of them) has a large vacuum expectation value which induces a mass for the inflaton of hilltop inflation, driving its vacuum expectation value to zero. The initial conditions for hilltop inflation are then generated when the preinflaton approaches a ``critical point'' where its stabilising effect ends and the fields enter a ``diffusion region'' in which quantum fluctuations have to be included.  We have calculated and discussed the resulting initial conditions  in section~\\ref{sec:ics} and found that there are two cases:  When the mass of the preinflaton lies above a certain threshold, initial conditions are generated such that the inflationary dynamics resembles almost entirely single field hilltop inflation. Single field hilltop inflation acts as an attractor solution where the initial value of the preinflaton becomes negligible when hilltop inflation starts.  The second case happens when the mass of the preinflaton is below the above-mentioned threshold. Then, the motion of the preinflaton during hilltop inflation can no longer be neglected and inflation has to be analysed as a true two field inflation model, using e.g.\\ the $\\delta N$ formalism. We discussed this in section~\\ref{sec:numericalResults} based on a numerical analysis, where we also highlighted the differences between these two cases.  As mentioned above, in addition to providing the right initial conditions for hilltop inflation, ``matter field preinflation'' has the attractive additional feature of opening up a decay channel for the inflaton into SM fields which can lead to an efficient reheating of the universe (with high enough but not too high reheat temperature). This can help to avoid unwanted decays of the inflaton into hidden sector fields and may even allow to explain the baryon asymmetry of the universe via the non-thermal leptogenesis mechanism.  As an example, in section~\\ref{sec:reheating} we calculated reheating and non-thermal leptogenesis for an example model where one of the right-handed sneutrinos acts as the preinflaton. We found that successful non-thermal leptogenesis is indeed possible and imposes a relation between the symmetry breaking scale and the mass of the right-handed (s)neutrino after inflation, as well as lower bounds on both quantities. In an explicit model, these would be interesting additional constraints on the model parameters.  We conclude that the proposed class of models of ``hilltop inflation with matter field preinflation'' is a promising framework for embedding inflation into particle physics theories."
    },
    "1402/1402.0999_arXiv.txt": {
        "abstract": "{This paper presents emission line intensities of CO and C$^{\\circ}$ from the compact circumnuclear disk in the center of NGC~5128 (Centaurus~A) obtained with the {\\it Herschel Space Observatory} in the 400-1000 GHz range as well as previously unpublished measurements obtained with the ground-based observatories {\\it SEST}, {\\it JCMT} and {\\it APEX} in the 90-800 GHz range.  The results show that the Cen~A center has an emission ladder of CO transitions quite different from those of either star-burst galaxies or (Seyfert) AGNs.  In addition, the neutral carbon ([CI]) emission lines from the Cen~A center are much stronger relative to the adjacent CO lines than in any other galaxy.  The CO surface brightness of the compact circumnuclear disk (CND) is significantly higher than that of the much more extended thin disk (ETD) in the same line of sight.  LVG analysis of the CO line profiles decomposed into the constituent contributions show that the ETD is relatively cool and of low excitation, wheres the brighter CND is hotter and more highly excited.  Our PDR/XDR models suggest that most of the CND gas is relatively cool (temperatures 25 K - 80 K) and not very dense ($\\approx\\,300\\,\\cc$) if it is primarily heated by UV photons. A small fraction of the gas in both the CND and the ETD has a much higher density (typically $30 000\\,\\cc$).  A more highly excited, high-density phase is present in the CND, either in the form of an extreme PDR or more likely in the form of an XDR.  Such a phase does not occur in the part of the ETD sampled.  We have determined, for the first time, the molecular mass parameters of the CND. The total gas mass of the CND is $M_{CND}\\,=\\,8.4\\times10^{7}\\,\\Msun$, uncertain by a factor of two.  The CO-$\\h2$ conversion factor ($X_{CND}$) is $4\\times10^{20}$ (K $\\kms$)$^{-1}$ also within a factor of two.} ",
        "introduction": "\\begin{table} \\scriptsize \\caption[]{Log of Herschel observations} \\begin{center} \\begin{tabular}{ccccr} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Instru- & Transi-       & OBSID      & Date  & Integr. \\\\ ment    & tion          &            & Y-M-D & (sec)  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\it HIFI} & CO $J$=5-4 C    & 1342200931 & 2010-07-19 & 1863 \\\\ {\\it HIFI} & CO $J$=5-4 NW   & 1342200932 & 2010-07-19 &   86 \\\\ {\\it HIFI} & CO $J$=5-4 SE   & 1342200933 & 2010-07-19 &   86 \\\\ {\\it HIFI} & 13CO $J$=5-4 C  & 1342201090 & 2010-07-21 &  679 \\\\ {\\it HIFI} & 13CO $J$=5-4 NW & 1342201093 & 2010-07-21 &  353 \\\\ {\\it HIFI} & 13CO $J$=5-4 SE & 1342201095 & 2010-07-21 &  353 \\\\ {\\it HIFI} & C18O $J$=5-4 C  & 1342201091 & 2010-07-21 &  112 \\\\ {\\it HIFI} & CO $J$=6-5 C    & 1342200982 & 2010-07-19 &   82 \\\\ {\\it HIFI} & CO $J$=6-5 NW   & 1342200983 & 2010-07-19 &   82 \\\\ {\\it HIFI} & CO $J$=6-5 SE   & 1342200984 & 2010-07-19 &   82 \\\\ {\\it HIFI} & CO $J$=7-6 C    & 1342201711 & 2010-07-30 & 1753 \\\\ {\\it HIFI} & CO $J$=7-6 NW   & 1342201714 & 2010-07-30 &   94 \\\\ {\\it HIFI} & CO $J$=7-6 SE   & 1342201715 & 2010-07-30 &   94 \\\\ {\\it HIFI} & CO $J$=8-7 C    & 1342201728 & 2010-07-30 &   92 \\\\ {\\it HIFI} & CO $J$=8-7 NW   & 1342201729 & 2010-07-30 &   92 \\\\ {\\it HIFI} & CO $J$=8-7 SE   & 1342201730 & 2010-07-30 &   92 \\\\ {\\it HIFI} & CO $J$=9-8 C    & 1342200948 & 2010-07-19 & 1896 \\\\ {\\it HIFI} & CO $J$=9-8 NW   & 1342200949 & 2010-07-19 &   70 \\\\ {\\it HIFI} & CO $J$=9-8 SE   & 1342200950 & 2010-07-19 &   70 \\\\ {\\it HIFI} & CO $J$=10-9 C   & 1342201108 & 2010-07-21 &  257 \\\\ {\\it HIFI} & CO $J$=10-9 NW  & 1342201109 & 2010-07-21 &  257 \\\\ {\\it HIFI} & CO $J$=10-9 SE  & 1342201110 & 2010-07-21 &  257 \\\\ {\\it HIFI} & CO $J$=13-12 C  & 1342201775 & 2010-07-30 & 1896 \\\\ {\\it HIFI} & CO $J$=13-12 NW & 1342201776 & 2010-07-30 &   70 \\\\ {\\it HIFI} & CO $J$=13-12 SE & 1342201778 & 2010-07-30 &   70 \\\\ {\\it HIFI} & [CI] $J$=1-0 C  & 1342201089 & 2010-07-21 &  353 \\\\ {\\it HIFI} & [CI] $J$=1-0 NW & 1342201092 & 2010-07-21 &  112 \\\\ {\\it HIFI} & [CI] $J$=1-0 SE & 1342201094 & 2010-07-21 &  112 \\\\ {\\it HIFI} & [CI] $J$=2-1 C  & 1342201712 & 2010-07-30 &   88 \\\\ {\\it HIFI} & [CI] $J$=2-1 NW & 1342201713 & 2010-07-30 &   88 \\\\ {\\it HIFI} & [CI] $J$=2-1 SE & 1342201716 & 2010-07-19 & 1884 \\\\ {\\it HIFI} & [CII] C         & 1342213717 & 2011-02-04 & 8970 \\\\ {\\it HIFI} & [CII] NW        & 1342201643 & 2010-07-28 &  936 \\\\ {\\it HIFI} & [CII] SE        & 1342201644 & 2010-07-28 &  936 \\\\ {\\it HIFI} & [NII]           & 1342201778 & 2010-07-31 &  833 \\\\ {\\it SPIRE}& SSPEC           & 1342204037 & 2010-08-23 & 5041 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{center} \\label{herschellog} \\end{table}  \\nobreak Disks of dense dust and gas deeply embedded in the stellar body of a giant elliptical galaxy are the more easily identified remnants of smaller gas-rich galaxies that have fallen in. They are a transient phenomenon: on varying time-scales, the gas will be consumed by accretion onto a central black hole, expulsion by jets emanating from the nucleus, and by the in-situ formation of new stars rejuvenating the host galaxy stellar population. All three processes are currently active in NGC~5128, the host of the Fanaroff-Riley class I (FR I) radio source Centaurus A (Cen A).  \\nobreak The properties of these embedded disks are probably different from those of the interstellar medium (ISM) in spiral or star-burst galaxies (see e.g. the review by Henkel $\\&$ Wiklind 1997). Kiloparsec-sized embedded thin disks are directly exposed to the intense but UV-poor combined radiation from all stars in the host elliptical.  There is little differential rotation, and the lack of shear may locally favour the formation of massive groups of luminous stars in the disk affecting their surrounding ISM. On much smaller sub-kiloparsec scales, the dynamics of a denser circumnuclear disk may be only loosely related to those of the larger kiloparsec-sized extended disk. Instead, in such a circumnuclear disk both dynamics and excitation expected to be tied more closely to the properties of the super-massive black hole in the nucleus.  \\nobreak The emission lines from carbon and carbon monoxide are key to understanding the properties of the ISM in galaxies, as they provide almost all the cooling of the dense neutral gas.  Warm and tenuous gas is traced by ionised carbon (\\cii\\,), warm and dense gas by neutral carbon (\\ci\\,) and carbon monoxide (CO), and cold and dense gas by (CO) and its isotopologues. Although the two neutral carbon \\ci\\ lines at 492 and 809 GHz and CO lines with rest frequencies up to about 1000 GHz can be measured from the ground (albeit with increasing difficulty), the ionised carbon \\cii\\ fine-structure line at 1.9 THz requires a platform above most or all of the atmosphere.  \\nobreak NGC~5128 is the nearest ($D$ = 3.84 Mpc) giant elliptical (review by Israel 1998). Its central black hole is still accreting in the aftermath of a merger with a medium-sized late-type galaxy a few hundred million years ago (Graham, 1979; Struve \\etal 2010). The remnant ISM of the merged galaxy has redistributed itself into a warped, thin disk ('the extended thin disk' - hereafter called the ETD - which shows in projection as a dark band crossing the optical image of the galaxy (Dufour et al. 1979, Nicholson \\etal 1992). It has comparable amounts of atomic (\\hi) and molecular H$_{2}$ gas, and its total mass is uncertainly estimated at $1.5\\times10^{9}$ M$_{\\odot}$, about two per cent of the enclosed dynamical mass.  On significantly smaller scales, the nuclear black hole of $5\\times10^{7}$ M$_{\\odot}$ (Cappellari \\etal 2009) is obscured by a compact circumnuclear disk (CND) that feeds the black hole (Israel \\etal 1990).  \\nobreak Unlike the ETD, the compact CND is oriented at right angles to the luminous radio and X-ray jets emanating from the Centaurus A nucleus. This suggests that the CND and the jets are somehow connected, but it is not yet clear how that may be.  The CND is a well-defined entity with a diameter of $20''$ (400 pc)).  It has been seen in FIR continuum (Hawarden \\etal 1993) and $^{12}$CO line emission (Espada \\etal 2009), but was not yet studied in detail.  In this paper, we combine single-dish CO and [CI] observations of the Cen A nucleus obtained from the ground over many years, and from space. Although observations of the Cen~A nucleus distinguishing CND and ETD in the lower ($J_{upper}\\,\\leq\\,3$) transitions of CO are found in a few earlier papers (Israel \\etal 1990; Israel 1992; Rydbeck \\etal 1993), all data in this paper have much higher signal-to-noise ratios and are presented here for the first time. ",
        "conclusions": "\\begin{enumerate} \\item We have measured $\\co$ fluxes from the central region of NGC~5128 (Centaurus A) in transitions up to $J$=12-11, as well as the fluxes from the submillimeter \\ci and \\cii lines. In addition, we have presented high-S/N velocity-resolved profiles of $\\co$ emission in the $J$=1-0 through $J$=5-4 and of $\\thirco$ emission in the $J$=1-0 through $J$=3-2 transitions. \\item $\\co$ line luminosities, normalized to a $22''$ beam (410 pc) increase up to $J$=5-4, and then steadily drop to low values. This drop is more pronounced than that seen in Seyfert AGNs, and much more pronounced than seen in star-burst galaxies. \\item Both \\ci\\ lines are more luminous than the adjacent ($J$=4-3, and $J$=7-6) $\\co$ lines. This behaviour is not seen in other (star-burst or AGN) galaxies, and is thus unique to the Cen~A center. The \\ci\\ line intensities and their ratio to CO line intensities are inconsistent with standard PDR model predictions. \\item Detailed analysis shows that up to $J$=6-5, about a third of the $\\co$ flux observed in the normalised $22''$ beam is contributed by the extended thin disk (ETD) embedded in NGC~5128, and unrelated to the compact circumnuclear disk (CND) just contained within the normalised beam. The CND contains a high proportion of very dense gas, in contrast to the ETD. At transitions above $J$=6-5, the $\\co$ emission is completely dominated by the CND contribution. \\item We have decomposed the observed CO spectral line ladder into individual (total) CND and (representative) ETD CO ladders. The ETD ladder peaks in the $J$=4-3/$J$=5-4 transition and then drops rapidly.  The CND ladder peaks in the $J$=6-5 and $J$=7-6 transitions before dropping more slowly. \\item LVG and PDR/XDR model analysis of the CND $\\co$ fluxes shows that most of the molecular gas mass resides in a relatively cool ($T_{kin}$ = 25 - 80 K), not very dense ($n_{\\h2}\\,\\approx\\,300\\,\\cc$ gas phase if the CND is heated exclusively by UV photons (PDR). \\item A small fraction of the gas in the CND is more highly excited and has much higher densities (typically $3\\times10^{4}\\,\\cc$). \\item In the CND but not in the ETD a third, more highly excited, high-density phase must also be present, either in the form of an extreme PDR or in the form of an XDR. \\item The CND has a total gas mass $M_{CND}\\,=\\,8.4\\times10^{7}\\,\\Msun$ (uncertain by a factor of two) which is about $10\\%$ of the mass of the much larger ETD. \\item The CO-$\\h2$ conversion factor ($X_{CND}$ is $4\\times10^{20}\\,\\kkms$ (uncertain by a factor of two), about twice the local Milky Way factor $X_{SN}$. \\end{enumerate}"
    },
    "1402/1402.0950_arXiv.txt": {
        "abstract": "{} { We study the influence of relevant quantities, including the density of dark energy (DE), to the predicted Hubble outflow around a system of galaxies.  In particular, we are interested in the difference between two models: 1) The standard $\\Lambda$CDM model, with the everywhere constant DE density, and 2) the \"Swiss cheese model\", where the universe is as old as the standard model, but the DE density is zero on short scales, including the environment of the system. } { We calculate the current predicted outflow patterns of dwarf galaxies around the Local Group-like system, using different values for the mass of the group, the local dark energy density, and the time of ejection of the dwarf galaxies, treated as test particles. These results are compared with the observed Hubble flow around the Local Group. } {The predicted distance-velocity relations around galaxy groups are not alone very sensitive indicators of the dark energy density, due to the obsevational scatter and the uncertainties caused by the used mass of the group and  a range in the ejection times. In general, the Local Group outflow data are in agreement with the local DE density being equal to the global one, if the mass is  about $4 \\times 10^{12} M_{\\odot}$; a lower mass $\\la 2 \\times 10^{12} M_{\\odot}$ could suggest a zero local DE density. The dependence of the inferred DE density on the mass is a handicap in this and other common dynamical methods. This emphasizes the need to use different approaches together, for constraining the local dark energy density. } {} ",
        "introduction": "Peirani \\& de Freitas Pacheco (2008) and Chernin et al. (2009) made calculations on the present-epoch Hubble relation around galaxy groups, without and with local dark energy (DE) and pointed out differences expected in these two cases. For instance, with dark energy, the zero-velocity distance $R_0$ will be longer than without dark energy, for a fixed group mass. Teerikorpi \\& Chernin (2010) further considered, within the Einstein-Straus vacuole containing different amounts of dark energy, the expected present-day local Hubble relation and its link to the global Hubble law.  Here we extend the previous works to include a range of relevant parameters (mass, ejection time) in order to reach better understanding on the factors influencing the use of the local Hubble flow as a measure of the dark energy. We concentrate in the cases where the central mass (point mass or binary), keeps constant, in order to see clearly the basic predictions and problems which would also go over into any perhaps more realistic and still uncertain models (such as having a changing mass).  Along with the standard $\\Lambda$CDM model, with its constant dark energy density on all scales, it is relevant consider the \"Swiss cheese model\", where the universe has the same age as the standard model, but the DE density is zero on short scales. This could correspond to the case where dark energy (or analogous effects) operates on large scales only. ",
        "conclusions": "Our calculations show clear differences between the cases of $\\Lambda$CDM and no local dark energy, in the sense expected. In practice, the difference may be difficult to detect.  First, the observed distance--velocity relation is rather scattered, roughly covering a similar region in the $r-V$ diagram as the theoretical curves. Second, the independently known masses of the galaxy groups (even our Local Group) are uncertain. Thirdly, the evolution of a galaxy group contains uncertain elements, including the ejection time.  We saw that the Hubble flow around the Local Group to be approximately consistent with the condition $\\rho_{\\rm loc} = \\rho_{\\Lambda}$ requires a rather high mass $\\approx 4 \\times 10^{12} M_{\\odot}$. The usually cited $\\approx 2 \\times 10^{12} M_{\\odot}$ would suggest about zero local DE density. Indeed, one point to emphasize is the interdependence between the assumed mass and the inferred DE density. This occurs here and also in other common dynamical methods such as the virial theorem and the Kahn-Woltjer method (Chernin et al.  2009).  In fact, attempts to probe dark energy with nearby outflows were first made (e.g., Chernin et al. 2006; Teerikorpi et al. 2008) by using the zero-gravity radius (Eq.2). This method is also affected by the \"mass--dark energy degeneracy\". If the possible range of $R_{\\rm ZG}$ is somehow estimated from the velocity-distance diagram\\footnote{According to our calculations the Hubble relation of the outflow does not show any clear signature of the zero-gravity distance. However, Teerikorpi \\& Chernin 2010 showed that if the outflow reaches the global Hubble rate at a distance $R_2$ from the group, the zero-gravity radius is expected to lie at $\\la R_2/1.5$. For the Local Group this is roughly where the outflow begins to be seen. } and the mass $M$ of the group is independently known or assumed, the DE density may be estimated in the outflow region: \\be \\frac{\\rho_{\\rm loc}}{\\rho_{\\Lambda}} = (\\frac{M}{2 \\times 10^{12} M_{\\odot}}) (\\frac{1.3 {\\rm Mpc}}{R_{\\rm ZG}})^3.\\ee Here the DE density is directly proportional to the mass. Hence Eq.5 may give useful constraints on the DE density for methods where the  $\\rho_{\\rm loc}$ vs. $M$ dependence is shallower, such as the  modified Kahn-Woltjer method, which results in a non-horizontal elongated area of admissible values in the local DE density-mass diagram. Thus Chernin et al. (2009) derived using the relative velocity and distance of our Galaxy and the M31 galaxy together with the outflow data the local DE density as $\\rho_{\\mathrm loc}/\\rho_{\\Lambda} = 0.8-3.8$. Using a similar approach, Byrd et al. (2014) revised this range to $\\rho_{\\mathrm loc}/\\rho_{\\Lambda} \\approx 0.5-2.5$.    To conclude,   $\\bullet$ The outflows around galaxy groups, though theoretically clearly influenced by local dark energy, are not very sensitive indicators of the exact value of the dark energy density in view of the obsevational scatter in the distance-velocity diagrams, as already noted by Peirani \\& de Freitas Pacheco (2008), and the uncertainties caused by the adopted mass of the group, a range in the assumed ejection times, and the binary nature of the group. Especially the uncertainties in mass and time are harmful.  $\\bullet$   In general, the Local Group outflow data are in agreement with the local DE density being equal to the global one, if the mass is  about $4 \\times 10^{12} M_{\\odot}$. A lower mass $\\la 2 \\times 10^{12} M_{\\odot}$ could suggest a zero local DE density.  $\\bullet$ Theoretically, local dynamics on galaxy group (a few Mpc) scales significantly differs for the cases of local DE = global DE or = 0 (e.g., Chernin et al. 2009, Byrd et al. 2012), though any  single method (e.g. local outflows, the Kahn-Woltjer method, the virial theorem) have large uncertainties. Hence, it is desirable to consider different methods together in order to lower systematic and random errors.  Up to now, no local evidence has appeared contradicting the concept of the universal cosmological constant as adopted in the $\\Lambda$CDM cosmology. Of course, such evidence if found would be very challenging for the standard cosmology (see, e.g., Byrd et al. 2012).  {\\it \\bf Acknowledgements}  We are grateful for discussions with G. Byrd, A.C. Chernin, and M. Valtonen.  \\appendix"
    },
    "1402/1402.2352_arXiv.txt": {
        "abstract": "{\\surfsup} is a joint {\\Spitzer} and {\\hst} Exploration Science program using 10 galaxy clusters as cosmic telescopes to study $z\\gtrsim 7$ galaxies at intrinsically lower luminosities, enabled by gravitational lensing, than blank field surveys of the same exposure time. Our main goal is to measure stellar masses and ages of these galaxies, which are the most likely sources of the ionizing photons that drive reionization. Accurate knowledge of the star formation density and star formation history at this epoch is necessary to determine whether these galaxies indeed reionized the universe. Determination of the stellar masses and ages requires measuring rest frame optical light, which only {\\Spitzer} can probe for sources at $z\\gtrsim 7$, for a large enough sample of typical galaxies. Our program consists of 550 hours of Spitzer/IRAC imaging covering 10 galaxy clusters with very well-known mass distributions, making them extremely precise cosmic telescopes. We combine our data with archival observations to obtain mosaics with $\\sim 30$ hours exposure time in both ${\\chone}$ and ${\\chtwo}$ in the central $4\\arcmin \\times 4\\arcmin$ field and $\\sim 15$ hours in the flanking fields. This results in 3-$\\sigma$ sensitivity limits of $\\sim26.6$ and $\\sim 26.2$\\,AB magnitudes for the central field in the IRAC 3.6 and ${\\chtwo}$ bands, respectively. To illustrate the survey strategy and characteristics we introduce the sample, present the details of the data reduction and demonstrate that these data are sufficient for in-depth studies of $z\\gtrsim 7$ sources (using a $z=9.5$ galaxy behind {\\macsb} as an example). For the first cluster of the survey (the Bullet Cluster) we have released all high-level data mosaics and IRAC empirical PSF models. In the future we plan to release these data products for the entire survey. ",
        "introduction": "\\label{sec:intro}  {\\surfsup} ({\\Spitzer} UltRa Faint SUrvey Program: Cluster Lensing and {\\Spitzer} Extreme Imaging Reaching Out to $z\\gtrsim 7$, \\#90009 PI Brada\\v{c}, co-PI Schrabback) is a joint {\\Spitzer} and {\\hst} Exploration Science program. It was designed to image 10 galaxy cluster fields to extreme depths with {\\Spitzer} ${\\chone}$ and ${\\chtwo}$ bands for 550 hours total. It also includes 13 prime and 13 parallel orbits of {\\hst} time for one of the clusters which did not have deep WFC3-IR and optical {\\hst} data (RCS2-2327.4$-$0204; the rest of the targets have {\\hst} data available).  Together with the archival data, each field has been or will be imaged with {\\Spitzer} for $>100\\mbox{ks}$ (28 hours) per band. Such depths have only been achieved previously with {\\Spitzer} observations of the Ultra Deep Field UDF \\citep{labbe12, vgonzalez10, labbe10}, GOODS (40 hours per field, see below and e.g., \\citealp{oesch13}) and CANDLES through the S-CANDELS program (P.I. G. Fazio; five CANDELS fields to 50 hours depth with IRAC). In the near future, SPLASH Survey (Spitzer Large Area Survey with Hyper-Suprime-Cam, PI Capak, \\#90042) will provide 2475h of Spitzer observing over two $1.8\\deg^2$ fields (COSMOS and SXDS); delivering depths of $\\sim 10$ hours per pointing. Compared to these studies, {\\surfsup} has the advantage of studying intrinsically lower luminosities, enabled by gravitational lensing, than blank field surveys of the same exposure time and has been designed to address the two main science goals described below.  \\subsection{Star formation at $z\\gtrsim 7$} The epoch of reionization marked the end of the so-called ``Dark Ages'' and signified the transformation of the universe from opaque to transparent. Yet the details of this important transition period are still poorly understood.  A compelling but most likely overly simplistic suggestion is that star-forming galaxies at $z\\gtrsim 7$ are solely responsible for reionization. The ability of sources to reionize the universe depends in part upon their co-moving star formation rate density $\\rho_{\\rm SFR}$ and star formation history at high redshift (for reviews, see \\citealp{fan06, robertson10, loeb12}). The advent of Wide Field Camera 3 (WFC3) on {\\hst} detects these galaxies at rest-frame UV wavelengths (e.g., \\citealp{ellis13, bouwens13,schenker13}), while {\\Spitzer} observations allow us to trace the rest-frame continuum emission redward of $4000\\mbox{\\AA}$ (e.g., \\citealp{labbe12}). Rest-frame UV and rest-frame $4000\\mbox{\\AA}$ data trace two basic properties of stellar populations; the instantaneous star formation rate (SFR) dominated by younger stars and the integrated history of the older population, respectively \\citep{madau99}. The stellar masses allow us to determine the SFR density at $z \\gtrsim 7$, which can be compared to the SFR density needed for these sources to reionize the universe (for certain choices of escape and clumping factors, \\citealp{madau99, robertson13, stark13}).  {\\surfsup} has the advantage that by using deep observations of 10 independent sight lines sample variance is reduced compared to e.g., UDF. Clusters of galaxies, when used as cosmic telescopes, allow us to probe deeper due to high magnification and {\\surfsup} targets are among the largest galaxy clusters known and were chosen for their extreme lensing strength. This program therefore allows us to push the intrinsic luminosity limits further than the UDF and study representative galaxies at $z\\sim 7$ and 8. For example, clusters that are part of this survey have typical magnifications of $\\mu \\gtrsim 5$, which effectively increases the exposure time by $\\sim \\mu^2$. For the $z\\sim 9.5$ galaxy reported below the {\\it intrinsic} (corrected for lensing) measured magnitudes in IRAC are $28.6^{+0.9}_{-0.8}$ in $\\chone$ and $27.9^{+0.6}_{-0.4}$ in $\\chtwo$, compared to 5-$\\sigma$ limiting magnitudes reported by \\citep{oesch13} in GOODS-N of 27.0 and 26.7, respectively.   One concern, however, when using gravitational lensing is that lensing magnification decreases the effective observing field (as it ``enlarges'' sources and their separations on the sky).  This loss in sky area is more than compensated for by the steep luminosity function (effective slope $>2$) at the magnitudes that we probe \\citep{bouwens12, bradley12}. A second concern is that we need to know the magnification (including errors) of our cosmic telescopes to convert the observed number counts and stellar masses into their intrinsic values. As shown by \\citet{bradac09}, the magnification of well-studied clusters, needed for such conversion, can be constrained using information on distortion and shifts of the background sources to sufficient accuracy. In summary, (1) in the regimes where the luminosity function is steep (effective slope $>2$, which is true at the magnitudes that we probe) number counts are increased compared to observations in a blank field (i.e., many somewhat fainter galaxies become accessible because of the foreground lens), and (2) magnification errors amount to a smaller error than sample variance when determining the luminosity function at $z\\sim 7$. Another advantage of gravitational lensing is that lensed galaxies are often enlarged, easing identification (gravitational lensing magnifies solid angles while preserving colors and surface brightness).   The first demonstration of an established stellar population at high redshift ($z\\gtrsim 6$) was accomplished using {\\Spitzer} data of the strongly-lensed $z\\sim 6.8$ galaxy behind Abell 2218 \\citep{egami05, kneib04}. Detections at ${\\chone}$ and ${\\chtwo}$ allowed the construction of the galaxy's spectral energy distributions (SED) and measurement of the stellar properties. The SED has a significant rest-frame $4000\\mbox{\\AA}$ break and therefore indicates that a mature stellar population is already in place at such a high redshift \\citep{egami05}. These measurements were made possible due to large magnification factors ($\\sim 25$). When observing gravitationally magnified objects, {\\Spitzer}/IRAC imaging enables us to study stellar populations of the highest redshift galaxies (e.g., see \\citealp{zheng12} for a $z\\sim 9.5$ galaxy detected by {\\Spitzer}; \\citealp{smit13} for detections at $z = 6.6-7.0$).     Considerable investment has recently been made in observing galaxy clusters with HST. The Cluster Lensing And Supernova survey with Hubble (CLASH; \\citealp{postman12}) delivered observations of 25 clusters and {\\hst}-GO-11591 (PI Kneib) observed an additional 9 clusters.  Future high redshift exploration will be advanced by the {\\hst} Frontier Field HFF\\footnote{\\url{http://www.stsci.edu/hst/campaigns/frontier-fields/}} program, a program involving six deep fields centered on strong lensing galaxy clusters in parallel with six deep \u201cblank fields\u201d (PIs Mountain, Lotz). Very deep {\\Spitzer} data are an excellent complement to deep {\\hst} data, which CLASH does not provide. The typical integration times for CLASH clusters prior to {\\surfsup} range from $\\sim 3.5$ hours per IRAC band from the ICLASH program (\\#80168: PI Bouwens, \\citealp{bouwens12c}) to $\\sim 5$ hours per IRAC band from the Spitzer IRAC Lensing Survey program (\\#60034: PI Egami). {\\surfsup} provides the greater depth and coverage needed in 10 strong lensing clusters specifically chosen for their high lensing strength (see below, 2 of them are part of HFF). The {\\Spitzer} campaign covering the HFF will provide similar depth for at least additional 2 clusters. In summary, {\\Spitzer} plays a unique role in the investigation of stellar ages and masses of $z\\gtrsim 7$ galaxies. IRAC ${\\chone}$ and ${\\chtwo}$ observation probe rest-frame optical wavelengths ($\\sim 0.5\\mu\\mbox{m}$) which are the only available data redward of rest-frame $4000\\mbox{\\AA}$ for these sources and hence can probe presence of evolved stellar populations for a large number of distant sources.   \\subsection{Evolution of Stellar Mass Function in Galaxy Clusters} With the {\\surfsup} observations, we will also be able to probe the stellar mass function of $z\\sim 0.3-0.7$ members of our cluster sample to depths of $<10^8$ M$_{\\odot}$ (or 0.005 L$^*$) for an elliptical galaxy at the highest cluster redshift we probe ($z=0.7$). This depth far exceeds the current limits from studies of other high-redshift clusters (e.g., \\citealp{andreon06b, patel09, demarco10, vanderburg13}) and is comparable to the deepest observations of local clusters, such as the Coma cluster and the Shapley Supercluster \\citep{terlevich01,merluzzi10}. Furthermore, it is comparable to the state-of the art stellar mass surveys at low redshifts; e.g., the Galaxy And Mass Assembly (GAMA) survey \\citep{taylor11, baldry12}, which has limits of $M^{*}\\sim 0.5-1 \\times 10^8$ M$_{\\odot}$ at a median redshift of $z = 0.05$. Previous optical/near-IR observations of galaxy clusters at $z \\sim 0.8$ suggested a deficit of faint, red galaxies in the cluster red sequence (RS) as indicated by the color-magnitude diagram (CMD) and the RS luminosity function (e.g., \\citealp{delucia04,delucia07, tanaka05,rudnick09, gilbank10, lemaux12}). However, other authors find no such deficit (e.g., \\citealp{depropris13, crawford09,andreon06, andreon08}). Because rest-frame optical luminosities can be strongly affected by current star-formation, color--stellar mass plots can look substantially different than CMDs (e.g., \\citealp{lemaux12}). As a result, the stellar mass function and the processes governing its evolution are the most physical and accurate way to trace evolution in the cluster galaxy population. So far, there has been little observed evolution in the cluster stellar mass function; however, published results have only probed down to $\\sim 10^{10}$ M$_{\\odot}$ (e.g., \\citealp{bell04, demarco10, vulcani13}).    Because evolution is accelerated in overdense environments (e.g., \\citealp{tanaka08}), it is essential to probe to lower stellar mass limits in the cluster cores to get a complete picture of galaxy evolution in these regions. {\\surfsup} will achieve that by making a complete census of star forming cluster galaxies down to stellar masses of $10^8$ M$_{\\odot}$ (or 0.005 L$^*$). Combined with our optical and near-IR photometry, IRAC data yield precise stellar masses and their errors ($< 0.15$ dex) for a particular choice of an initial mass function (IMF; e.g.,\\citealp{rowan-robinson08, swindle11}). The primary systematic uncertainty is the unknown IMF; for example, changing it from the Chabrier IMF \\citep{chabrier03} to the Salpeter IMF \\citep{salpeter55} will lead to a shift of $\\sim 0.25$ dex in stellar mass \\citep{swindle11}. Without the IRAC data, the statistical errors in stellar mass would increase by a factor of two. In summary, IRAC observations allow us to estimate stellar masses for all of our observed galaxies, down to a stellar mass limit comparable to that reached in local clusters \\citep{terlevich01,merluzzi10}.   This paper describes the survey design, key science goals, and details of reducing the ultra deep {\\Spitzer} data. We show the power of {\\surfsup} to achieve the primary goal listed above by measuring stellar properties for a $z=9.5$ galaxy behind {\\macsb}. In \\citet{surfsup2} we present details of the photometry and measurements of the stellar masses and SFRs for $z\\sim 7$ galaxies behind the Bullet Cluster. The full analysis of all 10 clusters, which will allow us to answer the questions described above, will be presented in subsequent papers after the final data is taken. The paper is structured as follows. In Section~\\ref{sec:survey} we describe the {\\surfsup} program, in Section~\\ref{sec:datared} we present the data reduction steps. In Section~\\ref{sec:key} we present the main science goal of the survey.  We summarize our conclusions in Section~\\ref{sec:conclusions}. Throughout the paper we assume a $\\Lambda$CDM concordance cosmology with $\\Omega_{\\rm m}=0.27$, $\\Omega_{\\Lambda}=0.73$, and Hubble constant $H_0=73{\\rm\\ kms^{-1}\\:\\mbox{Mpc}^{-1}}$ \\citep{komatsu11, riess11}.  Coordinates are given for the epoch J2000.0, and magnitudes are in the AB system. ",
        "conclusions": "\\label{sec:conclusions} {\\surfsup} will produce a major advance in our understanding of the formation of the first galaxies, in particular regarding their star formation history and stellar properties. This program will enable us to probe smaller stellar masses ($\\sim 10^8 M_{\\odot}$) and specific star formation rates ($\\sim 10^{-9} \\mbox{ yr}^{-1}$) at the highest redshifts $z \\gtrsim 7$. If these high-redshift galaxies are responsible for reionization, they need to produce a sufficient number of Lyman-continuum photons in a sustained way. Once these galaxies are identified, the IGM-ionizing photon flux will be estimated from the star formation rate density, which will include contributions from  instantaneous star formation rate dominated by younger stars and the integrated rate given by the older population \\citep{robertson13}.   In this paper and in \\citet{surfsup2} we have demonstrated the importance of using IRAC data to estimate stellar masses, ages, and SFRs for $z\\gtrsim 7$ galaxies. In particular, we have shown that without IRAC data the stellar properties are not robustly determined. At $z \\sim 7$, the addition of IRAC photometry in SED fitting significantly reduces the biases in the estimated galaxy properties compared to using HST photometry alone  \\citep{surfsup2}. At $z \\sim 9$, the lack of IRAC photometry in SED fitting can even lead to an order-of-magnitude bias in stellar mass, SFR and age estimates. Hence, {\\surfsup} will contribute significantly to accurate measurements of the stellar mass properties for these galaxies and thus, help constrain the IGM-ionizing photon flux.   Not only do we have a limited knowledge of the earliest formation of galaxies, but our picture of galaxy formation at later times is also lacking many details. The magnifying power of galaxy clusters also allows us to explore otherwise unreachable populations at ``intermediate'' redshifts ($1<z<7$). We will be able to probe the conditions in typical low-mass, star-forming galaxies at an at an epoch when they are otherwise inaccessible. The magnified galaxies provide excellent targets for exploiting the unique capabilities of new facilities like ALMA and JWST. By studying galaxy clusters, {\\surfsup} will also enable measurements of the stellar mass function of $z\\sim 0.3-0.7$ galaxy cluster members. The survey reaches depths of $<10^8$ M$_{\\odot}$ (or $<0.005 L^*$) for an elliptical galaxy at $z = 0.7$ (our highest redshift clusters). The large FOV of {\\Spitzer} will allow us to study cluster members out to $R_{\\rm vir}$.   Finally, {\\surfsup} will be a resource for the broader community for the study of distant, magnified sources and IR properties of lower redshift galaxies. Data for 9 out of 10 clusters have been taken. We have made available high-level science products (mosaics and empirical PSF measurements) for the Bullet Cluster and we plan on releasing all the data in the near future."
    },
    "1402/1402.1398_arXiv.txt": {
        "abstract": "{The chemical evolution of water through the star formation process directly affects the initial conditions of planet formation. The water deuterium fractionation (HDO/H$_2$O abundance ratio) has traditionally been used to infer the amount of water brought to Earth by comets. Measuring this ratio in deeply-embedded low-mass protostars makes it possible to probe the critical stage when water is transported from clouds to disks in which icy bodies are formed.} {We aim to determine the HDO/H$_2$O abundance ratio in the warm gas in the inner 150 AU for three deeply-embedded low-mass protostars NGC~1333-IRAS~2A, IRAS~4A-NW, and IRAS~4B through high-resolution interferometric observations of isotopologues of water.} {We present sub-arcsecond resolution observations of the $3_{1,2}-2_{2,1}$ transition of HDO at 225.89672~GHz in combination with previous observations of the $3_{1,3}-2_{2,0}$ transition of H$_2^{18}$O at 203.40752~GHz from the Plateau de Bure Interferometer toward three low-mass protostars. The observations have similar angular resolution (0\\farcs7 - 1\\farcs3), probing scales $R\\lesssim150$~AU. In addition, observations of the $2_{1,1}-2_{1,2}$ transition of HDO at 241.561~GHz toward IRAS~2A are presented to constrain the excitation temperature. A direct and model independent HDO/H$_2$O abundance ratio is determined for each source and compared with HDO/H$_2$O ratios derived from spherically symmetric full radiative transfer models for two sources. } {From the two HDO lines observed toward IRAS~2A, the excitation temperature is found to be $T_\\mathrm{ex}=124\\pm60$~K. Assuming a similar excitation temperature for H$_2^{18}$O and all sources, the HDO/H$_2$O ratio is $7.4\\pm2.1\\times10^{-4}$ for IRAS~2A, $19.1\\pm5.4\\times10^{-4}$ for IRAS~4A-NW, and $5.9\\pm1.7\\times10^{-4}$ for IRAS~4B. The abundance ratios show only a weak dependence on the adopted excitation temperature. The abundances derived from the radiative transfer models agree with the direct determination of the HDO/H$_2$O abundance ratio for IRAS~16293-2422 within a factor of 2-3, and for IRAS~2A within a factor of 4; the difference is mainly due to optical depth effects in the HDO line.} {Our HDO/H$_2$O ratios for the inner regions (where $T>100$~K) of four young protostars are only a factor of 2 higher than those found for pristine, solar system comets. These small differences suggest that little processing of water occurs between the deeply embedded stage and the formation of planetesimals and comets. } ",
        "introduction": "Water is a very important molecule through all stages of star and planet formation. In dense clouds, it is one of the dominant reservoirs of oxygen, both as a gas at high temperatures or as ice in grain mantles in cold regions. It also aids in the sticking of grains outside the snow line, and is crucial for the emergence of life on Earth. This makes it vital to understand the chemical evolution of water during the different stages of star formation, and to identify the mechanism by which it was brought to the Earth and Earth-like planets.   One way to gain insight into the chemical evolution of water is to compare the level of deuterium fractionation in water (the HDO/H$_2$O ratio) at various stages of star and planet formation. In our own solar system, several measurements of the HDO/H$_2$O ratio\\footnote{$\\mathrm{D/H}=0.5\\times\\mathrm{HDO/H_2O}$} exist for long-period comets from the Oort cloud \\citep[][and references therein]{mumma11}. Up until recently the mean HDO/H$_2$O ratio for these objects was $6.4\\pm1.0\\times 10^{-4}$ \\citep{villanueva09,jehin09}, significantly above the value for the Earth's oceans $3.114\\pm 0.002\\times 10^{-4}$ \\citep[VSMOW\\footnote{Vienna Standard Mean Ocean Water},][and ref. therein]{delaeter03}. Both values are well above the protosolar nebula value of 2$\\times$D/H=$0.42\\pm0.08\\times10^{-4}$ \\citep{geiss98,lellouch01} indicating that deuterium fractionation of water has taken place. Because of the observed difference between the Earth and comets, it has been argued that only a small fraction ($\\le10\\%$) of the Earth's water could have its origin in comets \\citep{morbidelli00,drake05}.  The Jupiter family comet Hartley~2 observed with the \\herschel\\ Space Observatory was found to have a HDO/H$_2$O ratio of $3.2\\pm0.5\\times10^{-4}$, significantly lower than the value for Oort cloud comets, and comparable to VSMOW \\citep{hartogh11}. Jupiter family comets have relatively short periods ($\\lesssim20$~years), with orbits more or less confined to the ecliptic plane, and originate in the Kuiper belt \\citep{duncan04,mumma11}. Oort cloud comets have longer periods, orbits distributed isotropically, and originate in the Oort cloud. In 2012 \\citeauthor{bockelee12} observed the Oort cloud comet Garradd with \\herschel\\ and deduced a HDO/H$_2$O ratio of $4.12\\pm0.44\\times10^{-4}$. This is significantly lower than previous measurements of Oort cloud comets. Recently, \\citet{lis13} measured the HDO/H$_2$O ratio in the Jupiter family comet Honda-Mrkos-Pajdusakov (HMP) using \\herschel\\ to $<4.0\\times10^{-4}$. These latest measurements indicate a certain range in the water deuterium fractionation in comets, possibly dependent on the precise formation location and subsequent migration, in which Jupiter family comets no longer have a distinctly lower HDO/H$_2$O ratio than Oort cloud comets. Simulations suggest that the deuterium fractionation in the protosolar nebula disk increases with distance from the forming sun \\citep{hersant01,kavelaars11} -- but whether the observed ratios are consistent with these simulations is still unclear.  Measurements of the water deuterium fractionation toward protostars provide a different perspective. By determining the HDO/H$_2$O ratio of material that enters protoplanetary disks, it is possible to set the initial level of deuterium fractionation in comet forming zones as input for any subsequent evolutionary models. So far, determinations of HDO/H$_2$O have given a wide range of results. \\citet{parise03} used ground-based infrared observations of the stretching bands of OH and OD in water ice in the cold outer parts of protostellar envelopes and found \\textit{upper limits} ranging from 0.5\\% to 2\\% for the HDO/H$_2$O ratios. Studies using gaseous HDO lines differ in the interpretation, with cold HDO/H$_2$O ratios derived from single-dish data ranging from cometary values \\citep[$2\\times10^{-4}$,][]{stark04}, to a few $\\times10^{-2}$ \\citep{parise05, liu11, coutens12}.  In the inner warm ($T>100$ K) region, HDO/H$_2$O ratios $\\geq$1\\% have also been inferred based on analysis of many HDO and water lines ($E_\\mathrm{u}$ up to 450~K) \\citep{liu11,coutens12}. In contrast, \\citet{jorgensen10b} derived a HDO/H$_2$O ratio toward the Class~0 protostar NGC~1333-IRAS~4B of $<6.4\\times10^{-4}$ from the interferometric Submillimeter Array (SMA) and Plateau de Bure (PdBI) observations, up to two orders of magnitude lower.   There are two main issues with determining the water abundances in the warm gas in the inner few hundred au. First, the relatively large beam sizes ($>10$\\arcsec) of single-dish telescopes are more sensitive to extended emission than the small $\\sim 1\\arcsec$ scales on which the warm water is found. Second, many observed water lines suffer from high optical depths, even for isotopologues, making it difficult to infer reliable column densities and abundances for H$_2$O in the warm ($>100$~K) gas. Indeed, \\citet{visser13} find that the $3_{1,2}-3_{0,3}$ lines ($E_u=249$~K) of H$_2^{16}$O, H$_2^{18}$O and even H$_2^{17}$O observed by the HIFI instrument \\citep{graauw10} aboard \\herschel\\ probably show optically thick emission on $R\\sim100$~AU scales in the well-studied Class 0 protostar NGC1333-IRAS~2A. Taking the optical depth properly into account lowers the HDO/H$_2$O ratio for IRAS~2A from $\\ge0.01$ to $1\\times10^{-3}$ \\citep{visser13}.  Recent high-resolution interferometric ground-based observations have made it possible to estimate the abundances of both H$_2$O and HDO in the warm gas on small scales. \\citet{jorgensen10a} targeted the H$_2^{18}$O $3_{1,3}-2_{2,0}$ line at 203 GHz with PdBI, which has a much lower Einstein $A$ coefficient than the lines observed with \\herschel\\ and thus suffers less from optical depth effects. Using SMA data of the HDO $3_{1,2}-2_{2,1}$ transition at 225 GHz, the above mentioned HDO/H$_2$O limit of $<6.4\\times10^{-4}$ was found toward NGC~1333-IRAS~4B \\citep{jorgensen10b}. \\citet{coutens13} obtained values of $4 - 30\\times10^{-4}$ for IRAS~4A and $1 - 37\\times10^{-4}$ for IRAS~4B using new \\herschel\\ data combined with the H$_2^{18}$O column densities found by \\citet{persson12} based on interferometric data.   \\citet{taquet13b} derived the warm HDO/H$_2$O ratios for IRAS~2A and IRAS~4A using the determinations of the H$_2^{18}$O column density from \\citet{persson12} and lower spatial- and spectral-resolution interferometric observations targeting the HDO $4_{2,2}-4_{2,3}$ transition (143.7~GHz, $E_\\mathrm{u}=316.2$~K) together with single dish \\citep{liu11} observations of HDO.  Using non-local thermal equilibrium (non-LTE) large velocity gradient (LVG) analysis, ratios in the range of $30-800\\times10^{-4}$ for IRAS~2A and $50-300\\times10^{-4}$ for IRAS~4A were obtained, depending on the assumed density, with the high-density ($n_\\mathrm{H_2}=1\\times10^8$~cm$^{-3}$) solution giving the lowest ratios.  Finally, using Atacama Large Millimeter/submillimeter Array (ALMA) and SMA observations, \\citet{persson13a} estimated the HDO/H$_2$O ratio of the warm water toward IRAS~16293-2422~A to $9.2\\pm2.6\\times10^{-4}$.  Most of these new values, summarized in Table~A.1, are significantly lower than the previous estimates of a few percentage points, indicating that the deuterium fractionation in warm water may not be as enhanced as previously thought.  In order to settle the question of the HDO/H$_2$O abundance ratios in warm gas, we have obtained high spatial resolution interferometric observations of HDO to complement previous high quality data of H$_2^{18}$O for three sources. The observations are sensitive to scales comparable to disks ($\\sim 1\\farcs3$, $R \\lesssim150$ AU), allowing us to determine the HDO/H$_2$O ratio in the innermost parts of young protostars where the disk is forming.  This paper presents observations of HDO \\trans{3}{1}{2}{2}{2}{1}\\ at 225.9~GHz toward the three well-known deeply-embedded low-mass protostars in the NGC~1333 region of the Perseus molecular cloud (\\citealt{jennings87,sandell91} --- IRAS~2A, IRAS~4A, and IRAS~4B in the following) at a distance of 250~pc \\citep{enoch06}. The three sources are ideal to compare because of their similar masses, luminosities, and location \\citep{jorgensen06}. \\citet{persson12} presented observations of the H$_2^{18}$O \\trans{3}{1}{3}{2}{2}{0}\\ transition at 203.4~GHz toward the same sources. With these two datasets we can deduce a HDO/H$_2$O ratio in the warm gas in the inner 150~AU radius using the same method and similar sensitivity and resolution as for IRAS~16293-2422, bringing the total number of low-mass protostars for which high quality data exist to four. In addition, observations of HDO $2_{1,1}-2_{1,2}$ at 241.6~GHz toward IRAS~2A were acquired to empirically constrain the excitation temperature.  The paper is laid out as follows. Section~\\ref{section:observations} describes the details of the various observations. Sections~\\ref{section:results} and \\ref{section:analysis} present the results and analysis -- including intensities, column densities, spectra, maps, and line radiative transfer modeling of the observed emission. In Sect.~\\ref{section:discussion} we discuss the deduced HDO/H$_2$O ratio and compare it to other studies of protostars and with solar system objects. ",
        "conclusions": "\\label{section:discussion}   All deduced HDO/H$_2$O ratios are consistent with our previous low estimates for two deeply-embedded protostars \\citep{persson13a,jorgensen10b}. The consistently low HDO/H$_2$O ratios for all four sources suggests that the chemical evolution of water for these sources is similar.  Our high-resolution observations make it possible to focus on the compact warm gas.  The sources show some extended emission in HDO and/or H$_2^{18}$O that would be included in any lower-resolution (single-dish) observations. Another key to our conclusions is that the observed lines are optically thin or at most marginally optically thick.  Because of their higher Einstein A coefficients, the $3_{1,2}-3_{0,3}$ lines of H$_2^{16}$O and H$_2^{18}$O observed with \\herschel\\ rapidly become optically thick, each at slightly different radii complicating the analysis \\citep{visser13}. Unless the optical depth is properly treated, the HDO/H$_2$O ratios deduced using \\herschel\\ data are likely to be overestimated. The ratio estimated for IRAS~2A by \\citeauthor{visser13} accounting for the optical thickness of the \\herschel\\ lines gives a HDO/H$_2$O ratio of $1\\times10^{-3}$, in agreement with our estimates.  The recently deduced warm HDO/H$_2$O ratios for IRAS~2A ($3\\times10^{-3}$) and IRAS~4A ($5\\times10^{-3}$) at the relevant densities ($n_\\mathrm{H_2}=1\\times10^{8}$~cm$^{-3}$) by \\citet{taquet13b} are factors of 3--5 higher than our estimates. With a lower spectral ($4$\\ vs. $0.1$~\\kms) and spatial (2\\arcsec\\ vs. 0.8\\arcsec) resolution than that of our H$_2^{18}$O data, the HDO observations used by \\citeauthor{taquet13b} probe larger scales, thus are affected by extended emissions as observed in our observations at scales $\\gtrsim500$~AU (with higher HDO/H$_2$O ratios as a probable result). Furthermore, the upper energy levels of the HDO and H$_2^{18}$O transitions differ significantly, $E_\\mathrm{u}=319$~K for the HDO $4_{2,2} - 4_{2,3}$ transition vs.\\ 204~K for the H$_2^{18}$O $3_{1,3} - 2_{2,0}$. The transitions could originate in slightly different regions depending on the source structure.  \\citet{coutens13} used \\herschel\\ and several ground-based single dish telescopes to observe lines of HDO toward IRAS~4A and IRAS~4B. From 1D radiative transfer models and using the H$_2^{18}$O column density by \\citet{persson12}, they estimated HDO/H$_2$O ratios in the inner warm regions of IRAS~4A ($4.0-30.0\\times10^{-4}$) and IRAS~4B ($1.0-37.0\\pm10^{-4}$). To reproduce the absorption seen in some of the spectral lines, a cold foreground absorbing layer is introduced. The results for the warm HDO/H$_2$O ratio show a wide range but are consistent with our findings, namely that the ratio is similar for the two protostars, and that it is at most a few $\\times10^{-3}$.  The radiative transfer modeling highlights the uncertainties involved in the derivation of the HDO/H$_2$O ratio. However, the ratios deduced with the optically thin LTE model-independent (i.e., not employing physical models and radiative transfer calculations to interpret the emission) method used in this paper and in \\citet{persson13a} is direct, and agrees with the full modeling within a factor of $3-4$. Until the physical structure on small scales can be constrained, running full radiative transfer models can not yield a more accurate HDO/H$_2$O ratio in the warm gas. However, spherically symmetric radiative transfer modeling can give indications on non-LTE effects and is a valid approximation for the outer envelope \\citep[e.g.,][]{mottram13}. Furthermore, the modeling also confirms that the observed transitions of the water isotopologues are not masing in these environments.  The emerging picture of the water deuterium fractionation in warm inner regions of deeply-embedded low-mass protostars is therefore that of a consistent and relatively low ratio. While the protostars in this study are all located in the same star forming cloud, NGC~1333, IRAS~16293-2422 in \\citet{persson13a} is located in a different cloud, $\\rho$~Ophiuchus. Its similar HDO/H$_2$O ratio indicates that a different birth environment does not necessarily imply a different ratio. To confirm this, HDO/H$_2$O ratios determination toward more sources are needed.   The infall timescales for the early, deeply-embedded stages are generally too short for any significant gas-phase chemical processing of water to take place in the warm $>$100 K gas before the material enters the disk \\citep{schoier02}. Therefore the gaseous water observed here probably originates by sublimation of ice formed earlier in the evolution.   In Fig.~\\ref{figure:dhfractionation} the deuterium fractionation of water for all four low-mass protostars (IRAS~2A, IRAS~4A-NW, IRAS~4B and IRAS~16293-2422) is presented together with those found in comets, Earth, interstellar medium (ISM) and protosolar values. The small differences in the amount of deuteration seen toward the inner regions of protostars and the comets in our solar system suggests that only small amounts of processing of water takes place between the deeply embedded stages and the formation of comets and planetesimals. If water was delivered to Earth by comets and small solar system bodies during the early evolution of the Sun, this would be the same water that was formed directly on the grain surface in the cold early stages of the formation of our Sun.  \\begin{figure*}[ht] \\centering \\includegraphics{fraction} \\caption{Values of the HDO/H$_2$O ratio, for different objects. For references, see Table~\\ref{table:ratios}. Error bars for protostars reflect those obtained using the optically thin LTE approximation.  } \\label{figure:dhfractionation} \\end{figure*}"
    },
    "1402/1402.1021_arXiv.txt": {
        "abstract": "{Many chemically peculiar (CP) stars, especially the magnetic CP2 stars, show a flux depression at 5200\\AA. The $\\Delta a$ photometric system takes advantage of this characteristic to detect these objects in an efficient way. In addition, it is capable of finding metal-weak, emission-type, and shell-type objects of the upper main sequence.} {To compare available observations and to detect new peculiar objects, we used a spectrophotometric catalogue consisting of 1159 stars. From this catalogue, we selected 1067 objects to synthesize three different $a$ indices to find the most efficient one for further observations. In addition, we extended the analysis to stars cooler than F5.  } {We employed classical $\\Delta a$ photometry described by Maitzen, using simulated filter curves, the spectrophotometric $\\Delta a$ index by Adelman, and a modified index.} {Even though the accuracy of the spectrophotometry used for this investigation is significantly lower than the photometric $\\Delta a$ measurements, we are able to confirm peculiarity for most of the known CP2 stars above a certain limit of $\\Delta a$. We investigated 631 stars hotter than spectral type F5 to find additional that are not yet identified peculiar objects. We find that for very low mass stars (M0), the $a$ index is independent of the colour (effective temperature).} {The $\\Delta a$ photometric system is very closely correlated with the effective temperature over a wide range of the main sequence. It is able to detect any kind of peculiarity connected to the 5200\\AA\\ region. Especially for low-mass stars, this opens up a new possibility of detecting peculiar objects in an efficient way.} ",
        "introduction": "Chemically peculiar (CP hereafter) stars of the upper main sequence differ in their abundances of heavy and rare-earth elements in their photosphere from normal-type objects of the same effective temperature and luminosity range. \\citet{Pres74} divided this group into four subgroups according to their astrophysical characteristics. Since then several investigations about their local environment and stellar parameters were performed. \\citet{Babc47} discovered a global dipolar magnetic field in the star 78 Virginis, which was soon followed by the detection of several other similar stars. Therefore the correlation of stellar magnetic field strengths with astrophysical processes such as diffusion and meridional circulation can be very well studied in this stellar group \\citep{Glag13}. Another important often studied aspect is the cause of the peculiar (surface) abundances. The observed phenomenon can be explained by either classical diffusion of chemical elements depending on the balance between gravitational pull and uplift by the radiation field through absorption in spectral lines \\citep{Stif12} or selective accretion from the interstellar medium via the stellar magnetic field \\citep{Havn71}. While diffusion seems to be appropriate for both magnetic and non-magnetic stars to explain spectral peculiarity, it is not yet clear to which extent the interaction with the interstellar medium via accretion and transport of angular momentum may modify the effects of diffusion and break the stellar rotational velocities during the stars' main sequence life time.  Using spectrophotometry, \\citet{Koda69} was the first to notice broad-band flux depressions at 4100, 5200 and 6300\\AA\\ during his investigation of the magnetic CP (CP2) star HD 221568. Photometrically, the main depressions (4100 and 5200\\AA) of another CP2 star were later also found by \\citet{Mait72} when they investigated the spectrum-variable HD 125248. This resulted in the development of the $\\Delta a$ photometric system \\citep{Mait76}, which makes use of the most efficient depression at 5200\\AA . The flux depression itself is a combined contribution of Si, Cr and Fe for strongly overabundant surface metallicities \\citep{Khan07}. The upper main sequence is unique for CP stars. No detailed observational analysis of this region for cooler objects has been performed so far.  In this paper, we present a study that was conducted to detect CP stars on the basis of spectrophotometric data for which we employed three different methods that use the flux depression at 5200\\AA. In addition to the original method, the $a'$ system developed by \\citet{Adel79} and a newly modified method were investigated. Furthermore, we investigated the capability of these systems of finding new, previously undetected, peculiar stars including metal-weak and emission-type objects of the upper main sequence. For the first time, the low-mass regime was searched for the behaviour of the $a$ indices in correlation with the effective temperature and luminosity.   \\begin{figure} \\begin{center} \\includegraphics[width=85mm]{aa22300f1.eps} \\caption{Regression of 437 normal stars with reference to the different $a$ indices.} \\label{regression} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=9cm]{aa22300f2.eps} \\caption{Peculiarity indices of 631 stars in the spectral-class range from O6 to F5. The filled circles denote stars that are detected as peculiar within at least one index.} \\label{3sigma} \\end{center} \\end{figure} ",
        "conclusions": "We used the spectrophotometric data of the stellar catalogue by \\citet{Khar88} to synthesize three different ``$a$ systems''. These data cover the complete spectral range from low- to high-mass objects. For the first time, we presented  $a$ indices for stars cooler than F-type. Excluding low-quality data and early-type supergiants, 1067 stars were used to synthesize different peculiarity indices. Our main results are \\begin{itemize} \\item Most of the known classical magnetic chemically peculiar stars were detected. \\item We presented a list of about 50 normal-type objects across the complete spectral range that were detected. \\item The most efficient $a$ system is very similar to that previously employed by \\citet{Mait80a}. \\item The normality line of the $a$ system correlates with the effective temperature up to a spectral type of M0. \\end{itemize} Our analysis showed that spectrophotometric data can be used for calculating synthetic $a$ indices and for detecting peculiar objects across the complete spectral range up to M0. As next steps, we will observe cool-type objects to establish the corresponding correlations. In addition, we will use classification resolution spectra (for example from LAMOST) to search for chemically peculiar objects in a semi-automatic way."
    },
    "1402/1402.0607_arXiv.txt": {
        "abstract": "{We study the effects of kination and scalar-tensor pre-Big Bang Nucleosynthesis cosmologies on the non-resonant production of sterile neutrinos.  We show that if the peak of the production rate of sterile neutrinos occurs during a non-standard cosmological phase, the relic number density of sterile neutrinos could be reduced with respect to the number expected in the standard cosmology. Consequently, current bounds on active-sterile neutrino mixing derived from the relic energy density of sterile neutrinos could be greatly relaxed.  In particular, we show that the sterile neutrinos which could explain the anomalies found in short-baseline neutrino experiments are compatible with recent joint Planck upper limits on their contribution to the energy density of the Universe in a scalar-tensor or a low-reheating temperature pre-Big Bang Nucleosynthesis cosmology.}   \\begin{document} ",
        "introduction": "In the Standard Model of particle physics, there are three active neutrinos, $\\nu_\\alpha$, coupled to the W and Z weak gauge bosons through gauge couplings.  Extensions of the Standard Model can include one or more sterile neutrinos, $\\nu_s$, that are not directly coupled to the W and Z bosons, but are produced in the early Universe via mixing with the active neutrinos.  Barbieri and Dolgov~\\cite{Barbieri} and Dodelson and Widrow~\\cite{Dodelson:1993je} provided respectively the first analytical calculation of the production probability and the first analytic solution for the relic number density of sterile neutrinos produced non-resonantly in the early Universe, under the assumption of a negligible primordial lepton asymmetry (later corrected by a factor of 2~\\cite{oai:arXiv.org:hep-ph/0202122, Abazajian:2001nj}).  If there is a large lepton asymmetry, sterile neutrinos can be produced resonantly \\cite{Abazajian:2001nj,Shi:1998km} with a non-thermal spectrum which favors low energies.  One or two sterile neutrinos with mass of $\\mathcal{O}(1\\,{\\rm eV})$ have been proposed as a way to explain several anomalies observed in short-baseline neutrino experiments, which aim a beam of neutrinos at a nearby ($\\lesssim 1\\,{\\rm km}$ away) detector.  LSND \\cite{Aguilar:2001ty} and Mini-BooNE \\cite{AguilarArevalo:2010wv, Aguilar-Arevalo:2013pmq} observed an excess of electron (anti-)neutrinos above the expected level. This excess could be explained by the oscillation of muon (anti-)neutrinos into one or two sterile (anti-)neutrinos, and the subsequent oscillation of these sterile (anti-) neutrinos into electron (anti-)neutrinos.  In addition to the LSND and Mini-BooNE results, radioactive source experiments meant to test the Gallium solar neutrino experiments SAGE and GALLEX have observed an event rate that is lower than expected.  This effect can be explained by electron neutrinos oscillating into sterile neutrinos \\cite{Giunti:2010zu,Acero:2007su}.  A recent re-evaluation of the neutrino flux emitted by nuclear reactors \\cite{Mueller:2011nm, Huber:2011wv} has led to increased predicted fluxes compared to previous calculations \\cite{Schreckenbach:1985ep,Hahn:1989zr,VonFeilitzsch:1982jw,Vogel:1980bk}.  Previous reactor experiments have not observed this flux, a result which can be explained by assuming $\\bar{\\nu_e}$ disappearance due to oscillations with a sterile neutrino. The sterile neutrinos that could explain these anomalies have a best fit active-sterile mixing angle ${\\rm sin}^22\\theta \\simeq 0.08$ (and ranging from $0.03$ to $0.4$ at the $95\\%$ confidence level)~\\cite{Kopp:2013vaa}.  From now on we will refer to sterile neutrinos with mass $m_{\\nu_s}=1\\,{\\rm eV}$ and mixing angle ${\\rm sin}^2(2\\theta)=0.08$ as the ``LSND $\\nu_s.$''  Recently, the Planck satellite experiment has provided joint constraints on the effective number of relativistic degrees of freedom at recombination, $N_{eff}$, and the effective mass of a sterile neutrino, $m_{\\nu_s}^{eff}$, in different models for active and sterile neutrinos.  $N_{eff}$ is defined as \\begin{equation} N_{eff}=3.046+\\frac{\\rho_{\\nu_s}}{\\rho_{\\nu_\\alpha}}, \\label{neff} \\end{equation} where $3.046$ is the expected value of $N_{eff}$ in the absence of sterile neutrinos \\cite{Mangano:2005cc}.  $\\rho_{\\nu_s}/\\rho_{\\nu_\\alpha}$ is the ratio of the energy density of relativistic sterile neutrinos, $\\rho_{\\nu_s}=2\\int d^3p/(2\\pi)^3E_{\\nu_s}f_{\\nu_s}$ for each $\\nu_s$ species, and the energy density of a neutrino in thermal equilibrium is $\\rho_{\\nu_\\alpha}$.  The effective sterile neutrino mass is defined for non-relativistic sterile neutrinos as \\begin{equation} m_{\\nu_s}^{eff} = \\frac{\\sum  n_{\\nu_s}m_{\\nu_s}}{n_{\\nu_\\alpha}}, \\label{meff} \\end{equation} where the sum is over sterile neutrino species, and $n_{\\nu_\\alpha}$ is the number density of an active neutrino.  For relativistic sterile neutrinos in thermal equilibrium with the active neutrinos, each $\\nu_s$ species would have the same relic density as one of the active neutrinos, thus $\\rho_{\\nu_s}/\\rho_{\\nu_\\alpha}$ would in fact be the number of sterile neutrino species.  For a non-relativistic neutrino in thermal equilibrium before decoupling, its relic number density would be the same as the number density of an active neutrino, $n_{\\nu_\\alpha}$, and its relic energy density would be $\\rho_{\\nu_s}=m_{\\nu_s}n_{\\nu_\\alpha}$.  Thus $m_{eff}$ would be the sum of the masses of all the sterile neutrino species.  Assuming the existence of two massless active neutrinos and a massive active neutrino with mass $m\\simeq 0.06\\,{\\rm eV}$ fixed by the atmospheric mass splitting, and a massive sterile neutrino with either a thermal distribution or a distribution proportional to that of the active neutrinos, and combining Planck, WMAP, Baryon Acoustic Oscillation, and high multipole CMB data, Planck finds the $95\\%$ limits \\cite{Ade:2013zuv} \\begin{equation} N_{eff}<3.80 \\label{nval} \\end{equation} and for sterile neutrinos with mass $m_{\\nu_s}<10\\,{\\rm eV}$ \\begin{equation} m_{\\nu_s}^{eff}<0.42\\,{\\rm eV}. \\label{mval} \\end{equation} The conditions of these bounds apply to the sterile neutrino distributions we consider in this paper.  The measurement of $N_{eff}$ provides a bound on the mixing angle of sterile neutrinos that are relativistic at the time of recombination, i.e. with mass $m_{\\nu_s} \\lesssim 1\\,{\\rm eV}$, while sterile neutrinos with higher masses up to $10\\,{\\rm eV}$ are constrained by the measurement of $m_{\\nu_s}^{eff}$.  Refs.~\\cite{Mirizzi:2013kva} and \\cite{Valentino:2013wha} have applied these bounds to a 3 active + 1 sterile neutrino model, and find the sterile neutrino needed to explain short-baseline neutrino experiment anomalies is highly disfavored.  Indeed, Ref.~\\cite{Mirizzi:2013kva} finds that this sterile neutrino is excluded at more than $4\\sigma$ by the Planck results.  One way to relax these bounds is to have a reduced number of sterile neutrinos produced in the early Universe.  Many possible scenarios have been considered, including models with a  non-standard cosmology, such as a low reheating temperature\\cite{Gelmini:2004ah, Gelmini:2008fq}, a constant $w\\neq 1$ dark energy equation of state \\cite{Hannestad:2005gj,Hamann:2011ge,Giusarma:2011zq}, a time-varying dark energy component \\cite{Giusarma:2011zq}, and interactions between the dark matter and dark energy sectors \\cite{Giusarma:2011zq,Kristiansen:2009yx,Gavela:2009cy}, or non-standard particle physics such as models with a large neutrino asymmetry \\cite{Foot:1995bm,Chu:2006ua,Saviano:2013ktj}, and interactions between sterile neutrinos and non-Standard Model particles \\cite{Bento:2001xi, Dasgupta:2013zpn,Hannestad:2013ana}.  Heavier sterile neutrinos with mass $m_{\\nu_s}\\gtrsim {\\rm keV}$ could account for the whole of the dark matter.  In addition, if such sterile neutrinos carry away a sizable fraction of the energy emitted in a supernova explosion, asymmetric emission of the sterile neutrinos due to the presence of a strong magnetic field could explain the very large velocities of pulsars \\cite{Fuller:2003gy}.  The mass and mixing of sterile neutrinos that make up all of the dark matter is subject to constraints from X-ray measurements, phase space density arguments, and analysis of the Lyman-$\\alpha$ forest data.  If the keV mass sterile neutrinos are produced non-resonantly, they are strongly disfavored as the sole dark matter component \\cite{Boyarsky:2008ju}.  If these sterile neutrinos are not the primary dark matter component, however, these bounds are weakened \\cite{Palazzo:2007gz}, and a new region of parameter space could be opened, as long as the relic energy density of the keV mass sterile neutrinos does not exceed the energy density of dark matter.  This paper considers the effects of kination and scalar-tensor pre-Big Bang Nucleosynthesis (BBN) cosmologies on sterile neutrinos.  In both of these cosmologies, the Hubble expansion rate is greater than in the standard cosmology, and so the production of sterile neutrinos in the early Universe is suppressed.  For comparison we also include the results found in Ref.~\\cite{Gelmini:2004ah} for sterile neutrinos evolving in a low reheating temperature cosmology. ",
        "conclusions": "In this paper we presented the effects of some non-standard pre-BBN cosmologies on the production of sterile neutrinos.  If the peak of the production rate of sterile neutrinos occurs during the non-standard cosmological phase, we have shown that the number density of sterile neutrinos produced non-resonantly could be much smaller than in the standard cosmology. Thus, current bounds on the active-sterile neutrino mixing derived from cosmological bounds on the energy density of sterile neutrinos can be greatly relaxed if the Universe developed differently than the standard radiation dominated cosmology assumes.  We find that if the Universe evolves in a scalar-tensor or low reheating temperature cosmology prior to BBN, a sterile neutrino that could explain the anomalies observed in short-baseline neutrino experiments would be allowed by the Planck joint bounds in Eq.~\\ref{nval} and \\ref{mval}, although it would be disfavored if the Universe evolved according to the standard or kination cosmologies.  In general, if the expansion rate during the non-standard cosmological phase depends on temperature as $T^n$, with $n\\geq 0,$ and goes smoothly into the standard cosmology, the LSND $\\nu_s$ production would never be in equilibrium only if $n>4.7.$  If the non-standard expansion rate has a quick jump at the transition into the standard cosmology, and has a very small $n,$ the condition is instead that the expansion rate in the non-standard cosmology at the time of the peak interaction rate must be $8.9$ times that of the standard cosmology, and to evade the Planck joint bounds, this number should be $21.$  We also find that the kination, scalar-tensor, and low reheating temperature cosmologies open up a new region of phase space for sterile neutrinos of mass $m_{\\nu_s}\\gtrsim {\\rm keV}$ that do not make up all of the dark matter.  The experimental discovery of such a neutrino could, therefore, be evidence of a non-standard pre-BBN cosmology."
    },
    "1402/1402.3268_arXiv.txt": {
        "abstract": "We derive physical properties of 10 submillimeter galaxies located in the CANDELS coverage of the GOODS-S field. The galaxies were first identified as submillimeter sources with the LABOCA bolometer and subsequently targeted for 870$\\mu$m continuum observation with ALMA. The high angular resolution of the ALMA imaging allows secure counterparts to be identified in the CANDELS multiband dataset. The CANDELS data provide deep photometric data from UV through near-infrared wavelengths. Using synthetic spectral energy distributions, we derive photometric redshifts, stellar masses, extinction, ages, and the star formation history. The redshift range is z$=$1.65-4.76, with two of the galaxies located at z$>$4. Two SMG counterparts have stellar masses 2-3 orders of magnitude lower than the rest. The remaining SMG counterparts have stellar masses around $1\\times10^{11}$\\,M$_{\\odot}$. The stellar population in the SMGs is typically older than the expected duration of the submillimeter phase, suggesting that the star formation history of submillimeter galaxies is more complex than a single burst. Non-parametric morphology indices suggest that the SMG counterparts are among the most asymmetric systems compared with galaxies of the same stellar mass and redshift. The HST images shows that 3 of the SMGs are associated with on-going mergers. The remaining counterparts are isolated. Estimating the dust and molecular gas mass from the submm fluxes, and comparing with our stellar masses shows that the gas mass fraction of SMGs is $\\sim$28\\% and that the final stellar mass is likely to be $\\sim(1-2)\\times10^{11}$\\,M$_{\\odot}$. ",
        "introduction": "Submillimeter galaxies (SMGs) are the most powerful starbursts known in the universe. With star formation rates, derived from their far-infrared luminosity, often exceeding several $10^{3}$\\,M$_{\\odot}$\\,yr$^{-1}$, SMGs are capable of building up a stellar mass of $\\sim10^{11}$\\,M$_{\\odot}$ in less than 100 Myr. This formation process of SMGs resembles the monolithic collapse scenario (Eggen et al. 1962), but might be driven by hierarchical processes that merges two gas rich galaxies. The role of accretion and minor mergers is currently unknown, but could play a role in the slightly less luminous SMGs. It has been proposed that the end result of the SMG process are the giant elliptical galaxies seen in the local universe (Blain et al. 2002).  Despite a substantial observational effort over the last 15 years, there are still unanswered questions concerning SMGs: the redshift distribution, the stellar mass content, what are the progenitors and descendants,  and finally, whether the SMGs play a significant role in the overall formation and evolution of galaxies or do they just represent the most extreme manifestation of a general star forming galaxy population?  The main hurdle in learning more about the SMG population has from the start been the difficulty in identifying optical and near-infrared counterparts. SMGs are initially identified using single dish bolometer arrays, with an angular resolution $\\ga10''$. By pushing interferometric radio continuum observations to levels where massive star formation is the main contributor to the continuum emission, and with sub-arcscond angular resolution, it was possible to identify the most likely counterparts of some SMGs and thereby obtain optical spectroscopic observations (e.g. Chapman et al. 2003; Aretxaga et al. 2007; Smolcic et al. 2012a; Hodge et al. 2013). This method, however, can at best only account for a limited fraction of all SMGs. A significant number of SMGs remain undetected even in the deepest radio continuum observations, possibly due to being at high redshift where the radio continuum falls below the detection limit, and some do not reveal optical emission and no spectroscopy can be done. An illustrative example of an elusive SMG is the brightest submillimeter source in the Hubble Deep Field North (Hughes et al. 1998). Despite a large observational effort, no optical/near-infrared counterpart has been identified. The redshift was only recently determined using molecular emission lines and found to be $z\\approx5.2$ (Walter et al. 2012). The counterpart remains undetected at other wavelengths. In addition to using radio continuum to identify optical/NIR counterparts, Spitzer mid-IR data has also been used (e.g. Pope et al. 2006; Biggs et al. 2011; Targett et al. 2013; Alberts et al. 2013).  To avoid the potential selection bias inherent in using the radio continuum to identify SMG counterparts, and to achieve a more direct association between the SMG and potential counterparts, it is better to directly target the millimeter and submillimeter emission using interferometric techniques with sub-arcsecond resolution. This method has successfully been used with the SMA, CARMA, Plateau de Bure and now also ALMA interferometers (e.g. Younger et al. 2008; Capak et al. 2011; Smolcic et al. 2012a, 2012b; Bothwell et al. 2013; Vieira et al. 2013; Hodge et al. 2013). These bias-free studies have found several SMGs at redshifts $z>4$, challenging both the previously derived redshift distribution of SMGs and the modeling of the formation of massive galaxies in the early universe (e.g. Baugh et al. 2005). These results suggests that a substantial part of the SMG population, $\\sim$30\\%, is located at $z \\ga 4$ (Capak et al. 2008; Coppin et al. 2009; Daddi et al. 2009a, 2009b; Cox et al. 2011; Smolcic et al. 2011, 2012a; Wagg et al. 2012; Swinbank et al. 2012), with at least five SMGs found at $z > 5$ (Capak et al. 2011; Combes et al. 2012; Walter et al. 2012; Vieira et al. 2013). A recent detection of a z$=$6.3 SMG was reported by Riechers et al. (2013), pushing the redshift limit of the SMG phenomenon to within $\\sim$1 Gyr from the Big Bang. Observations of CO emission lines in lensed SMGs detected with the South Pole Telescope (Vieira et al. 2013) suggests that the radio continuum identification of SMGs may have missed a significant fraction of $z\\ga4$ SMGs and there are indications that sources selected at longer wavelengths may have redshift distributions that peak at higher redshifts than 870$\\mu$m selected samples (e.g. Smolcic et al. 2012; Yun et al. 2012; Vieira et al. 2013).  Regardless of their redshift, SMGs tend to have a remarkably uniform molecular gas mass. Summarizing various CO detections from the literature, and adopting a uniform CO-H$_{2}$ conversion factor of 0.8\\,M$_{\\odot}$\\,(K\\,km\\,s$^{-1})^{-1}$, the typical molecular gas mass of SMGs is $\\sim3 \\times 10^{10}$\\,M$_{\\odot}$ with only a small dispersion. (e.g. Greve et al. 2005; Tacconi et al. 2008; Coppin et al. 2009; Bothwell et al. 2013 and references therein)  With the new ALMA interferometer it is possible to detect sub-millimeter continuum emission at a flux level of $\\sim1$\\,mJy with sub-arcsecond resolution in a matter of minutes. This will greatly simplify the identification of SMG counterparts and allow an unbiased study of the SMG population. The first results for SMGs observed with ALMA during the Early Science phase have recently been presented by Hodge et al. (2013) and Vieira et al. (2013). The ALMA observations of Hodge et al. (2013) targeted the 126 LESS\\footnote{LESS = LABOCA ECDF Submillimeter Survey; ECDF = Extended Chandra Deep Field. The sources observed with ALMA has been referred to as ALMA LESS, or ALESS. In the text we will for simpicity refer to the SMGs as LESS `NN' and use the designation of Hodge et al. (2013) for multiple sources within the LABOCA beam (see Weiss et al. 2009). The corresponding IAU designations are given in Table~\\ref{tab:sources}.} sources found with APEX/LABOCA (Weiss et al. 2009). Of the 126 LESS sources, 17 are located within the GOODS-S field (Giavalisco et al. 2004). The ALMA observations provide sub-arcsecond location on 9 of the 17 sources located within the CANDLES coverage of the GOODS-S field, in some cases detecting multiple submillimeter continuum sources within a single LABOCA footprint.  In this paper we use the coordinates for the SMGs observed by Hodge et al. (2013) to identify the optical/NIR counterparts of the SMGs within the CANDELS GOODS-S field. In total we identify and analyze the properties of 10 counterparts. We use the near-infrared CANDELS data (Grogin et al. 2011; Koekemoer et al. 2011; Guo et al. 2013) as well as optical BViz data from the GOODS project (Giavalisco et al. 2004), re-measured using PSF matching and the TFIT method (see Guo et al. 2013), and Spitzer/IRAC infrared data (Ashby et al. 2013). In a recent paper, Targett et al. (2013) also discussed properties of SMGs in the GOODS-S field, using HST/WFC3 data. In this case, the counterparts were identified using radio continuum as well as Spitzer/IRAC 8$\\mu$m data.  Our sample partly overlaps with that of Targett et al. (2013), with five sources in common. The focus of the two papers are different but complementary and we compare our results with that of Targett et al. when appropriate. We adopt H$_{0}=72$\\,km\\,s$^{-1}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout this paper. All magnitudes are in the AB system (Oke 1974). ",
        "conclusions": "\\label{discussion}  \\subsection{The SMG counterparts}\\label{identification} Submm interferometry represents the least biased method at our disposal to identify the counterparts of SMGs. In contrast to radio continuum as well as near- and mid-infrared emission techniques, submm interferometry targets the same emission region as initially detected using single-dish bolometer arrays. The interferometric method, however, relies on bolometer surveys to select the targets in the first place, and as such is not completely free of selection biases. Bolometer surveys result in flux limited samples at a wavelength of $\\sim$870$\\mu$m, and thus consist of the most luminous far-infrared sources, with dust temperatures that allows the detection of the Rayleigh-Jeans part of the continuum. Galaxies dominated by a hot dust component, for instance, may not be part of the sample (e.g. Magnelli et al. 2010).  The identification of the SMG counterparts can be fraught with ambiguities despite using high quality submm interferometry and the deepest and  highest angular resolution optical/near-infrared images available. In the SMG sample presented here, there are several instances where the submm emission region is clearly separated from the optical/near-infrared emission. In some cases, this simply reflects a real offset between a pre-existing stellar component and the region that is currently forming stars. Such offsets are present in some local interacting/merging galaxies, such as the Antennae Galaxy, but typically not for local ULIRGs, where optical/near-infrared, dust and molecular line emission originates in a very compact region in the center of the ULIRG (e.g. Sanders \\& Mirabel 1996).  Despite the ambiguities in defining the actual counterpart to SMGs, the use of  interferometers to directly associate the submm emission with optical/near-infrared counterpart(s) and/or nearby galaxies gives us a higher confidence in the redshift determination, regardless whether it is spectroscopic or photometric redshifts. At least two of the SMGs in the sample presented here are located at z$>$4, with one spectroscopically confirmed source. Given the small sample, we regard this as consistent with the high $z>4$ SMG fraction suggested in earlier surveys (e.g. Aravena et al. 2010; Smolcic et al. 2012a,b; Walter et al. 2012; Bothwell et al. 2013; Vieira et al. 2013). The reminder of our SMG sample falls in the `traditional' SMG redshift range $z=1.65 - 2.60$.   \\subsection{Stellar properties}\\label{stellar} The SMGs in our sample appear to have two types of counterparts in terms of stellar mass. One more massive, with an average stellar mass of $0.9 \\times 10^{11}$\\,M$_{\\odot}$, and one type with a stellar mass 2-3 orders of magnitude lower. The distribution of stellar masses is shown in Fig.~\\ref{fig:hist} and shows the remarkably small dispersion for the high mass galaxies. The low mass of LESS 10 and LESS 34 is not an artifact of the SED fitting process; their optical and near-infrared fluxes are intrinsically lower than the rest of the SMG sample while their submm properties are similar. It is tempting to view these two galaxies as being in an earlier evolutionary stage, where the main part of the stellar mass is about to be formed (see Sect~\\ref{nature}). Our estimates of the stellar masses for the high-mass SMGs are smaller than that of the galaxies believed to be the final descendants, giant elliptical galaxies, and are thus consistent with continued build-up of stellar mass through gas infall and/or merging.  One perplexing characteristic of the SMG stellar population is their relatively old age (up to 1-2 Gyr). The SFR's inferred from the far-infrared luminosity of a typical SMG can build up a $10^{11}$\\,M$_{\\odot}$ galaxy from scratch in $\\sim10^{8}$ years, and one would naively expect the SMGs to be dominated by a young stellar population. The presence of the old(ish) stellar population suggests that the star formation history of SMGs is more complex than envisioned. The parametrization of the star formation history used in the SED fitting process can strongly influence the results (e.g. Michalowski et al. 2012). The `standard' exponentially declining parametrization is a poor fit to a more complex star formation history. The delayed-$\\tau$ parametrization of the star formation history used in this paper is an improvement, allowing an active and on-going star formation period, including an instantaneous burst. It is, however, not able to model a multiple-burst star formation history. Multiple star forming periods are usually not used in SED fitting procedures because of the limited number of photometric data points. With too many parameters the number of degrees of freedom approaches zero and the goodness-of-fit eventually becomes meaningless. If the real population is composed of a young and an old stellar component, the simple, one-component star formation history will create a luminosity weighted average stellar population. The  ages implied by the SED fit will then underestimate the age of the evolved population while overestimating the young population. Hence, the stars in the SMG presented here could very well have a component that is even older. This could as well have an effect on the derived stellar masses. As the stellar population ages, the M/L ratio increases and for a fixed luminosity, the stellar mass increases. Targett et al. (2013) used a two-component stellar population and found stellar masses that on average are $\\sim$4 larger than the ones derived by us (see Table~\\ref{tab:targett}). It is currently not clear whether this discrepancy is due to the use of one versus two stellar components, or simply reflects differences in the SED fitting algorithms (e.g. Michalowski et al. 2012).  Apart from a stellar population consisting of distinct age groups, variation of other parameters across a galaxy can introduce a bias in the parameters derived in the SED fitting process. For instance, if a galaxy consists of regions with different amounts of dust extinction, the SED fit will reflect the averaging over the different regions. That this is affecting at least one of the SMGs in our sample is illustrated in Fig.~\\ref{fig:LESS18}, which shows the LESS 18 source in the ACS/F435W, F606W and the WFC3/F160W filters. The photometric redshift of LESS 18 is $z=1.81$ and the F435W filter probes restframe $\\sim1500$\\AA. The galaxy has a very different appearance in restframe UV compared to restframe V band (probed by the WFC3 F160W filter). The two bright regions seen in the F435W filter are either absent or very weak in the WFC3 F160W filter showing that these are young star forming regions with little dust obscuration, while the main component of the galaxy either contains an older stellar population and/or is more obscured. In this case the SED parameters are biased towards a younger and less massive stellar component than what would have been derived without the UV component. LESS 18 seems to be the only case in the present sample where there is a UV component that is not seen at longer wavelengths.  The high metallicities of the SMGs in our sample imply a prolonged star formation activity and are in accord with the older ages of the stars found through the SED fitting. There are, however, degeneracies between metallicity, age and dust extinction. If the ages are overestimated, it requires either the metallicity to be even higher or the dust extinction to be more severe in order for the stellar templates to fit the observed fluxes. Since both the implied metallicities and extinction values are already very high for the present SMG sample, it is difficult to include a significantly younger stellar population and still achieve a good fit to the observed data. This suggests that the young stars are likely to be heavily obscured and their contribution to the observed fluxes are small.   \\subsection{Morphology}\\label{morphdiscussion} One of the main questions regarding SMGs is if they are the result of the merging of two gas-rich galaxies, with the end result being a giant elliptical galaxy (e.g. Blain et al. 2002), or if the star formation activity is fueled through some other mean. In the merging scenario one would expect the SMG counterparts to show the presence of close neighbors and/or tidal features.  The SMGs presented here have some properties similar to local ULIRGs: large dust and gas-masses, a high star formation rate and a large far-infrared luminosity (Sect.~\\ref{gas}). However, only 3 of the SMGs in our sample show clear evidence for merging while 7 of them are either completely isolated (5) or have one or two neighbors (2) within 30\\,kpc but show no signs of interaction (Table~\\ref{tab:results}). This makes the SMGs different from local ULIRGs, which are almost exclusively major mergers (e.g. Sanders \\& Mirabel 1996). The local ULIRGs are powered by a combination of a circumnuclear starburst and AGN activity, fueled by a reservoir of molecular gas. In contrast, only one of the 3 SMGs showing evidence for merging contains an AGN (LESS 13). Of the remaining 7 non-merging SMGs, 5 show evidence of an AGN.  The fact that the majority of the SMGs appear to be isolated galaxies and only 30\\% are merger (and of the three mergers, only one is an equal mass merger), suggests that the SMGs are different from local ULIRGs, despite having similar properties in terms of star formation rate, dust and gas content. In fact, the type of isolated massively star forming systems implied by the results for our SMG's at high redshift are not present in the local universe. Hydrodynamical simulations of star forming galaxies at z$\\sim$2 (Dekel et al. 2009) suggests that about half of the SMGs with observed 870$\\mu$m fluxes $\\ga$5\\,mJy are fueled by mergers with a mass ratio $>$1:10. The other half is fueled by a smoother inflow that includes mergers with mass ratios $<$1:10. Given the small sample presented here, our results for the merger fraction and mass ratios are consistent with the results of Dekel et al. (2009). It is also possible that we are viewing the SMGs during different stages in their evolution, where the obvious merger systems represent an earlier phase. However, the obvious candidates for the early part of the SMG phase are LESS 10 and LESS 34, neither of which shows any evidence of merging. Also the gas, dust and SFR properties (Sect.~\\ref{gas}) are remarkably similar for the early candidates, the mergers and the `isolated' SMGs. High angular resolution molecular line observations might provide an answer to these questions by mapping the distribution and kinematical structure of the star forming gas.  Comparing the asymmetry parameter of the SMGs with galaxies of similar mass and redshift shows that the SMGs are on average more asymmetric, but does not distinguish whether they are mergers or just exhibits a clumpy structure. The dusty nature of SMGs could be, at least partially, responsible for the high asymmetry seen at restframe optical wavelengths. The fact that z$\\sim$2 Lyman break galaxies have properties similar to both SMGs and local ULIRGs in the M$_{20}$-Gini plot, while having very low asymmetry values, opens up the possibility that SMGs are disk like structures, similar to the LBGs, with a high asymmetry caused by either a clumpy structure or patchy dust distribution. Several studies using decomposition of the light profile of SMGs have shown them to be dominated by exponential disks (Targett et al. 2011, 2013) with an average Sersic index  $<n>\\approx1.4$.  Swinbank et al. (2010) find a slightly larger Sersic index $<n>\\approx 2$ for a sample of 25 SMGs observed with HST/NICMOS. These results suggest that the light profile of SMGs in general have a disk-like characteristic.  Kartaltepe et al. (2012) recently presented an analysis of the morphological parameters of a sample of 52 z$\\sim$2 ULIRGs using CANDELS data. The ULIRGs were selected based on their far-infrared properties as observed with Herschel. The morphology of the sample was classified through visual inspection. The merger fraction was found to be high, $\\sim$72\\% compared with $\\sim$32\\% for comparison sample with lower far-infrared luminosity. They also found that overall the fraction of disk-like systems is approximately equal to the fraction of spheroids, but that the fraction of disks decreases with increasing $L_{\\mathrm{FIR}}$. Hence, the low merger fraction found for the SMGs is perplexing. Based on a comparison with the Kartelpe et al. sample, the fraction of disk-like system should be very low for the SMGs, which is contradictory to the results found by Targett et al. (2013). High angular resolution kinematical data are needed to determine the morphology of the SMGs.   \\subsection{The nature of the SMG and their counterparts}\\label{nature} Our results regarding the properties of the SMGs and their counterparts show that in terms of their far-infrared, star formation and interstellar gas characteristics, they constitute a homogeneous sample, but in terms of their stellar properties they divide into two distinct groups; a low stellar mass group and a high stellar mass group. The latter has a typical stellar mass of $0.9\\times10^{11}$\\,M$_{\\odot}$ with a remarkably small dispersion (Fig.~\\ref{fig:hist}).  The possibility that the low stellar mass of LESS 10 and LESS 34 is due to $\\sim$99\\% of the stars being completely obscured by dust is unlikely as their far-infrared luminosity and gas properties are similar to the other SMGs. If less UV/optical light `leaks' out, it should be absorbed and re-radiated by dust, resulting in a higher far-infrared luminosity. The L$_{\\mathrm{FIR}}$ of LESS 10 and LESS 34 are typical for SMGs, with $(7-9)\\times10^{12}$\\,L$_{\\odot}$. In fact, LESS 10 and LESS 34 appear to be completely `normal' SMGs in their FIR, gas and SFR properties (see Table~\\ref{tab:sfr}). It is only the stellar mass of the designated counterpart that makes them special. The SED fit of LESS 10 and LESS 34 are greatly improved by removing a power-law continuum in the IR (Sect.~\\ref{photoz}). The magnitude of the implied power-law is strong enough that essentially all of the IRAC fluxes are removed. Performing the SED fit for LESS 10 and LESS 34 without the IRAC data, and without a continuum power-law, also improves the $\\chi^{2}$ values, while leaving the fit parameters essentially unchanged; in particular, the photometric redshift and stellar mass remain exactly the same. This suggests that the IRAC fluxes are anomalous and inspection of the SEDs (Fig.~\\ref{fig:sed_a}) shows that the IRAC fluxes of both LESS 10 and LESS 34 appear to be disjoint from the optical/NIR fluxes. In the case of LESS 10, there is a neighbor galaxy 0\\ffas8 away, which leads to the possibility of blending of the IRAC photometry. The SED fit of the neighbor, both with and without the IRAC fluxes, suggests that it is situated at a completely different redshift and, hence, not physically related to the designated counterpart of LESS 10. Nevertheless, we did a SED fit combining the fluxes for the two galaxies. The result is a very poor SED fit ($\\chi^{2}_{\\nu}\\sim$20), with a stellar mass $\\sim$5$\\times10^{9}$\\,M$_{\\odot}$, still $\\sim$20 times less massive than the high-mass SMG counterparts. In the case of LESS 34 there is no nearby galaxy that can cause blending of the IRAC fluxes.  The anomalous optical/NIR vs. IRAC fluxes of LESS 10 and LESS 34 could be explained if the submm and IRAC fluxes are associated with a background galaxy, not visible in any of the optical/NIR bands. While LESS 34 has a small offset between the optical/NIR and submm emission (although only 0\\ffas34), the submm emission of LESS 10 is very well centered on the designated counterpart (Fig.~\\ref{fig:smooth}). The only way of resolving this issue is to obtain spectroscopic data on molecular or atomic fine-structure lines.  Another possibility that has to be considered is that these two galaxies represents an early stage in the development of SMGs. The characteristic age of LESS 34 is the youngest of our sample (50\\,Myr), but LESS 10 is better fit by an older stellar population. The stellar age is, however, not a robust parameter and although it can be used as a diagnostic for a reasonably large sample, it should not be used for the analysis of individual galaxies. If all of the molecular gas in LESS 10 and LESS 34 is turned into stars, it would result in a stellar mass of $5\\times10^{10}$\\,M$_{\\odot}$ in $\\sim$40\\,Myr, the gas consumption time scale. If the star formation efficiency (SFE) is less than 100\\%, the resulting stellar mass would be lower by a factor of a few. Even with a high SFE, the resulting SMG counterpart would be at least 2-3 times less massive than the typical stellar mass of the other SMG counterparts. Furthermore, the high-mass SMGs are characterized by stars significantly older than 40\\,Myr. In addition, the high mass SMGs contain as much molecular gas and have as high L$_{\\mathrm{FIR}}$ as LESS 10 and LESS 34 and, hence do not seem to simply represent a later stage in the SMG episode. In order for LESS 10 and LESS 34 to transform into the high mass counterparts, they would have to undergo at least 2, possibly 3-4, episodes of similar star formation activity as they are currently experiencing. If these episodes occur with a time lag of $\\sim$0.5\\,Gyr, the resulting galaxy would resemble the high mass galaxies. While the current data and results do not prove such episodic SMG activity, it is consistent with such a scenario. If correct, a larger sample of SMGs would yield counterparts with stellar masses ranging from the $10^{8}$\\,M$_{\\odot}$, all the way up to $10^{11}$\\,M$_{\\odot}$, and the apparent dichotomy of stellar masses found here will disappear.  The presence of  low-mass counterparts to SMGs has not been seen in other studies of SMG stellar masses (cf. Wardlow et al. 2011; Michalowski et al. 2012; Simpson et al. 2013). Hainline et al. (2011), however, did find a few cases where the stellar mass of the counterparts approached what we find for LESS 10. Some of these studies used a simplified estimate of the stellar mass, based only on the H-band luminosity and a M/L ratio from stellar evolution models, and the identification of the counterparts was in some cases based on data with less angular resolution than used here. The combination of ALMA and HST data makes the identification of the SMG counterpart more precise than previously, and our stellar mass estimates have proven to be a robust outcome from the SED fitting.  The above discussion begets the question of what constitutes a `SMG episode'? Merger between gas-rich galaxies or, a more or less continuous accretion? The general assumption is that the most luminous SMGs are powered through a merger while less luminous SMGs are fueled through a more continuous accretion, possibly involving accretion of clumps and dwarf galaxies with a mass ratio $<$0.1 (cf. Finlator et al. 2006; Dekel et al. 2009; Dav\\'{e} et al. 2010). This is supported by the data presented here. The only verified equal mass merger in our sample is LESS 13, which is also the most far-infrared luminous SMG. The two other mergers (LESS 67 and LESS 79.2) have mass ratios of 1/5 and 1/20, respectively, and both have far-infrared luminosities below the average of our sample. In fact, LESS 79.2 is the least luminous object of our SMGs. If this type of un-equal mass mergers/accretion is fueling the star formation activity on time scales less than $\\sim$40\\,Myr, the gas is constantly replenished and stellar mass is built-up in a more or less continuous fashion. The gas accretion rate at z$\\ga$2 is estimated to be sufficient to sustain a SFR of several $10^{2}$\\,M$_{\\odot}$\\,yr$^{-1}$ (e.g. Dekel et al. 2009, 2013). At some point, however, the star formation is quenched and this seems to occur around $(1-2)\\times10^{11}$\\,M$_{\\odot}$ as more massive SMG counterparts are not found, or at least appear to be very rare."
    },
    "1402/1402.4162_arXiv.txt": {
        "abstract": "Zhu, Stone, and Rafikov (2012) found in 3D shearing box simulations a new form of planet-disk interaction that they attributed to a vertical buoyancy resonance in the disk. We describe an analytic linear model for this interaction. We adopt a simplified model involving azimuthal forcing that produces the resonance and permits an analytic description of its structure. We derive an analytic expression for the buoyancy torque and show that the vertical torque distribution agrees well with results of Athena simulations and a Fourier method for linear numerical calculations carried out with the same forcing. The buoyancy resonance differs from the classic Lindblad and corotation resonances in that the resonance lies along tilted planes. Its width depends on damping effects and is independent of the gas sound speed. The resonance does not excite propagating waves. At a given large azimuthal wavenumber $k_y > h^{-1}$ (for disk thickness $h$), the buoyancy resonance exerts a torque over a region that lies radially closer to the corotation radius  than the Lindblad resonance. Because the torque is localized to the region of excitation, it is potentially subject to the effects of nonlinear saturation. In addition, the torque can be reduced by the effects of radiative heat transfer between the resonant region and its surroundings. For each azimuthal wavenumber, the resonance establishes a large scale density wave pattern in a plane within the disk. ",
        "introduction": "Young planets can experience strong gravitational interactions with surrounding gas residing in a protoplanetary disk \\citep{GT80, LP86, W97, LI11, KN12}. These interactions can lead to structural changes in a disk through the launching of waves that can result in shocks and  gap formation. They can also lead to changes in the orbital properties of the planet, resulting in their radial migration. Such phenomena are caused by the resonant forcing of the gas by the planet. The two types of resonances that describe this interaction, the Lindblad and corotational, have been extensively explored.  Both of these resonances involve planar motions and will occur in a two-dimensional disk in which  the vertical dynamical effects (perpendicular to the disk orbital plane) are ignored. For planets on circular orbits, the corotation resonance lies at the orbit of the planet. The gas response at the corotation resonance is of the form of a trapped radially evanescent wave with a radial drop-off on a scale  of order the disk thickness $h$. The other form of resonance, the Lindblad resonance, occurs where a forcing frequency component due to the planet matches the epicyclic frequency of the gas. There are infinitely many of these resonances. In the absence of disk self-gravity, these resonances result in the launching of acoustic waves that transport energy and angular momentum away from the planet. The strongest Lindblad resonances lie close to the planet, but cannot occur closer than $\\sim  h$ radially from the orbit of the planet due to pressure effects.   When the disk vertical structure is taken into account, a richer set of waves can be excited at Lindblad resonances \\citep{LO98, Bate02}. These waves are of the form of r modes (rotation-dominated), g modes (vertical buoyancy-dominated),  p modes (pressure-dominated), and f modes (fundamental). For small or moderate azimuthal wavenumbers ($\\ll h^{-1}$), the Lindblad response is dominated by an f mode, while the  r-modes and g-modes are less strongly excited. In a vertically isothermal disk that undergoes adiabatic perturbations, as will be considered in this paper, the fundamental mode consists of two-dimensional planar motions.   By means of 3D shearing box simulations, Zhu et al (2012) have recently shown that a new form of planet-disk interaction can occur due to buoyancy resonances. For this type of resonance, the disk vertical motions play a critical role. At such a resonance, a planet forcing frequency component matches the free oscillation frequency of a vertically displaced fluid element that behaves adiabatically. Like the Lindblad case, there are infinitely many such resonances, one for each azimuthal wavenumber $k_y$. But since the free vertical oscillation frequency varies with  height above the disk midplane, these resonances do not occur at particular radii as in the Lindblad case. Instead they lie along tilted planes. These resonances radially extend much closer to the planet than Lindblad resonances. The one-sided (inside or outside corotation) torque that results from the buoyancy resonances was found to be comparable to, but smaller than, the usual Lindblad torque.  The goal of this paper is to explore the analytic properties of buoyancy resonances. The analytic approach provides a verification of the existence of this resonance and insight into its structure. We consider the linear response of an isothermal disk that undergoes adiabatic perturbations to a simplified form of forcing. As in Zhu et al (2012), we consider the gas to reside in a shearing box, as described in Section \\ref{beq}. We consider waves for which the azimuthal wavenumber is of order or greater than $h^{-1}$. These waves play an important role in determining the total torque caused by a planet. To make analytic progress, we consider  forcing that is purely azimuthal and is independent of radius and height in the disk. In Section \\ref{saf}, we analyze  a disk with constant vertical gravity and find separable solutions for the linear disk response. We obtain analytic expressions for the structure of the resonance and the torque distribution. We then consider in Section \\ref{safvg} the case of variable vertical gravity and determine the linear response numerically by means of a Fourier method in radius. We then show that the torque density in this case agrees well with the torque density obtained by an obvious extension of the  torque density expression in the constant gravity case. Section \\ref{disc} contains a discussion and Section \\ref{sum} contains the summary. ",
        "conclusions": "}  Buoyancy resonances have been previously analyzed in the context of high mass stars that are tidally perturbed by a companion \\citep{Z75, GN89}. A buoyancy  resonance occurs near the outer parts of the stellar convective core where the buoyancy frequency changes from low values in the core to high values in the radiative envelope. Within this region, the buoyancy frequency matches that of the tidal forcing. The response of the star is to launch buoyancy waves (g mode waves) that propagate towards the stellar surface where the waves damp and act to bring about spin synchronization with the binary orbit.  As we see in this paper, the adiabatic response of a disk to tidal forcing is evidently quite different. In the axisymmetric case of an isothermal disk,  buoyancy waves (g modes) can be excited at a Lindblad resonance \\citep{Bate02}. They propagate on the same side of the Lindblad resonance as the fundamental mode propagates which is not the region where the buoyancy resonance discussed here is found.  For a vertically truncated isothermal disk, they accumulate near the disk upper boundary. We do not find evidence for these waves in the numerical results presented here at small $|x| < h$.  Low frequency axisymmetric waves in the form of r  (rotation-dominated) modes exist and are affected by buoyancy that confines them near the disk midplane (see Fig. 11 of \\cite{LP93}). These modes can also be excited at a Lindblad resonance and propagate toward corotation ($x=0$), although they are only weakly excited there \\citep{Bate02}.   To understand the star-disk difference, consider  low frequency modes of given $k_y$  in the region where wave frequency $|\\omega|$ is smaller than epicyclic frequency $\\kappa$, which is the region where the buoyancy resonances described in this paper reside. In a region such that $k_y \\ll |k_x|$, the WKB dispersion relation for the disk pressure perturbation is \\begin{equation} k_x^2    =  \\frac{\\kappa^2 - \\omega^2}{\\omega^2 - N^2(z)} \\,  k_z^2, \\label{dr} \\end{equation} where $\\omega$ is the wave frequency and \\begin{equation} \\omega= 2 A k_y x \\end{equation} \\citep{VD90, G93}. This relation can be easily derived from equation (\\ref{peq}). We see then from equation (\\ref{dr}) that vertical propagation is possible ($k_z^2 >0$) only for low values of the buoyancy frequency, $N^2< \\omega^2$. For a disk with $N \\propto z$,  such a wave can propagate vertically only in the region near the disk midplane as an r mode.  Consider the stellar case. To see how waves can be launched towards  a stellar surface, we adapt equation (\\ref{dr}) to the case of a star by noticing that $k_z$ in a disk describes phase variations along the direction of the buoyancy gradient. In the case of a star, the buoyancy gradient is in the (spherical) radial direction, and so we identify $k_z$ with $k_r$. We identify  the square of the wavenumber perpendicular to the buoyancy gradient with $k_x^2  \\sim \\ell (\\ell+1)/r^2$ for the spherical harmonic of  order $\\ell$ that is associated with the tidal field. We disregard the stellar rotation and set $\\kappa=0$. The tidal forcing frequency due to the companion star is $\\omega$. We then have that \\begin{equation} k_r^2 = \\frac{\\ell (\\ell+1)}{r^2} \\frac{N^2- \\omega^2 }{\\omega^2}. \\label{drs} \\end{equation} Radial wave propagation is possible (i.e., $k_r^2 >0$) in the radiative outer layers of a star outside the buoyancy resonance  where $N^2(r) > \\omega^2$ (cf. equation (15) of \\cite{GN89} with $\\ell=2$). Waves can be launched at the resonance  where $N^2(r) \\simeq \\omega^2$ because the long wavelength of the wave $\\sim 1/k_r$ can match the spatial scale of the slowly varying tidal field. The effects of rotation modify the dispersion relation by changing $\\omega^2$ in the denominator on the right-hand side of equation (\\ref{drs}) to $\\omega^2-\\kappa^2$. Therefore, one major difference between the disk case investigated here and stellar case is due to the effects of rapid rotation, $\\kappa$, in the low frequency disk case.  Consider possibility that r modes are launched at buoyancy resonances in disks. Waves are launched at resonances if there is a strong overlap between the spatial form of the wave with that of the forcing. Typically that occurs at wave turning points where the long wavelength form of the wave matches the form of the relatively slowly varying potential. For example, from equation (\\ref{dr}), $k_x$ is small near a Lindblad resonance where $\\omega = \\pm \\kappa$. This locally long wavelength permits a strong coupling with the tidal potential that results in the excitation of a wave. Although vertical wave turning points for r modes lie on the buoyancy resonance plane ($k_z=0$ at $\\omega=N$ in equation (\\ref{dr})),  the radial wavenumbers $k_x$ are large, as follows from equation (58) of  \\cite{LP93} with dimensionless frequency $F= 2 A x k_y/\\Omega$ for $|F| \\ll 1$. (Note: \\cite{LP93} used the term g modes for what are r modes, see also \\cite{O98}.) The r mode becomes more confined vertically as it approaches corotation. Its vertical wavenumber $k_z$ at the midplane and radial wavenumber $k_x$ both grow to very large values near corotation ($x=0$). Therefore, the r mode structure may not generally  match the spatial form of the potential along the buoyancy resonance. This lack of matching may explain why r modes are  not excited at buoyancy resonances.   Differential equation (\\ref{peq1}), which we obtained for the  pressure perturbation near a buoyancy resonance in the constant gravity case, is consistent with dispersion relation (\\ref{dr}). In that case, $N$ is constant, $k_z^2 = -1/h^2$ (see equation (\\ref{psf})), and near the buoyancy resonance $N^2 - \\omega^2 \\propto x-x_{\\rm res} \\propto X$. We then have that $k_x^2 \\propto 1/X$ which covers the leading order terms on the left-hand side of equation (\\ref{peq1}), identifying $P''(X)$ with $-k_x^2 P(x)$. However, the WKB approximation does not provide much insight into the properties of a buoyancy resonance because it does not involve waves.  \\cite{LP93} determined the structure of the disk modes  by using what could be called the \"waveguide\" model. In that model, the amplitude variations in the $x$ direction are assumed to occur slowly compared to phase variations. For perturbations induced by the buoyancy resonance, that assumption does not apply. For the case of vertically varying gravity, amplitude variations along both the $x$ and $z$ directions are rapid and comparable."
    },
    "1402/1402.6217_arXiv.txt": {
        "abstract": "The detection of gravitational wave usually requires to match the measurement data with a large number of templates, which is computationally very expensive. Compressed sensing methods allow one to match the data with a small number of templates and interpolate the rest. However, the interpolation process is still computationally expensive. In this article, we designed a novel method that only requires to match the data with a few templates, yet without needing any interpolation process. The algorithm worked well for signals with relatively high SNRs. It also showed promise for low SNRs signals. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3604_arXiv.txt": {
        "abstract": "We present spectroscopic follow-up observations on two bright Ly$\\alpha$ emitter (LAE) candidates originally found by \\citet{KRUG12} at a redshift of $z \\sim 7.7$ using the Multi-Object Spectrometer for Infra-Red Exploration (MOSFIRE) at Keck. We rule out any line emission at the $>5\\sigma$ level for both objects, putting on solid ground a previous null result for one of the objects. The limits inferred from the non-detections rule out the previous claim of no or even reversed evolution between $5.7 < z < 7.7$ in the Ly$\\alpha$ luminosity function (LF) and suggest a drop in the Ly$\\alpha$ luminosity function consistent with that seen in Lyman Break galaxy (LBG) samples.  We model the redshift evolution of the LAE LF using the LBG UV continuum LF and the observed rest-frame equivalent width distribution. From the comparison of our empirical model with the observed LAE distribution, we estimate lower limits of the neutral hydrogen fraction to be 50~-~70\\% at $z~\\sim~7.7$. Together with this, we find a strong evolution in the Ly$\\alpha$ optical depth characterized by $(1+z)^{2.2\\pm0.5}$ beyond $z=6$ indicative of a strong evolution of the IGM. Finally, we extrapolate the LAE LF to $z\\sim9$ using our model and show that it is unlikely that large area surveys like UltraVISTA or Euclid pick up LAEs at this redshift assuming the current depths and area. ",
        "introduction": "Understanding when and how the universe re-ionized is fundamental to our understanding of how galaxies and large scale structure form and evolve and is sensitive to global cosmological parameters. In particular, the \\textit{fraction of neutral hydrogen}, $x_{HI}$, in the intergalactic medium (IGM) is closely tied to early galaxy formation because it is related to the gas accretion rate onto galaxies. From current measurements it is still unclear when re-ionization occurred and what the sources of re-ionizing radiation are.  The best of such current measurements come from cosmic microwave background (CMB) experiments and high-redshift quasar studies, with additional constraints from Lyman Break (LBG) and Ly$\\alpha$ emitting (LAE) galaxy studies. WMAP \\citep{LARSON11} and Planck \\citep{TAUBER10} place a $\\sim2-3\\sigma$ constraint on when re-ionization occurred, based on the optical depth to the CMB due to Thompson scattering of electrons. These data are usually fit by a quick re-ionization at $z\\sim10.5$, but are also fully consistent with a more gradual re-ionization with a tail ending at $z\\sim6-7$ \\citep{KOMATSU11,PLANCK13}. Direct measurements of the optical depth from quasars indicate that the universe is neutral up to $z\\sim7.1$, based on the highest redshift quasars known today \\citep{FAN06,MCGREER11,MORTLOCK11}. Furthermore, ultraviolet (UV) continuum measurements of LBGs between $z\\sim7-10$ \\citep{BOUWENS11,BRADLEY12,SCHENKER12,MCLURE13} suggest that galaxies have a difficult time re-ionizing the universe until later times unless the luminosity function is unusually steep at the faint end, or the continuum escape fraction is high \\citep{ROBERTSON13}.  The fraction of strong Ly$\\alpha$ emitters within LBG samples should give us a more direct, complementary, and unique measurement of $x_{HI}$ and therefore how quickly and when the universe is re-ionizing.   Fundamentally, Ly$\\alpha$ photons are scattered in areas where the IGM contains more neutral hydrogen, so the escape fraction of Ly$\\alpha$ photons is proportional to the volume of re-ionized hydrogen around the young galaxies. Hence the fraction of galaxies with strong  Ly$\\alpha$ emission is related to the neutral fraction of the IGM \\citep{HAIMAN99,MALHOTRA04,DIJKSTRA07,MALHOTRA06, DIJKSTRA10}.  However, it is important to note that this probe is also sensitive to the evolution of the interstellar medium (ISM) inside galaxies (like dust, see \\citet{BOUWENS12b,FINKELSTEIN12,MALLERY12}), so one must understand the effects of galaxy evolution to probe the IGM.  The Ly$\\alpha$ emission of LBG galaxies (selected using broad bands) is indicative of re-ionization ending at $z\\sim6-7$ and a neutral hydrogen fraction of $\\sim50$\\% at $z\\sim7$ \\citep{FONTANA10, STARK10,PENTERICCI11,ONO11,SCHENKER11,CARUANA13}. In particular the fraction of strong Ly$\\alpha$ emitters in LBG samples is found to rapidly drop beyond $z > 6.5$ over a range of $\\Delta z \\gtrsim 1$, a timescale of only $\\sim 200$ Myrs \\citep{STARK11,CURTISLAKE12,SCHENKER11}.  An alternative to LBG selection is the use of narrow-band (NB) filters to directly detect LAEs at specific redshifts (e.g., \\citet{MALHOTRA01,HU04} and references therein). This method allows one to directly map the Ly$\\alpha$ LFs as a function of redshift, which can then be compared to the LBG UV continuum LFs to estimate the neutral IGM fraction.  An overall change in the Ly$\\alpha$ LFs between $5.7 < z < 6.6$ has been firmly established by large samples of spectroscopically confirmed LAEs \\citep{OUCHI08,OUCHI10, HU10, KASHIKAWA11,MALHOTRA04}. But the source of this change could be either an evolution in the IGM or a change in the internal ISM of the galaxies.   The evolution of the Ly$\\alpha$ LF based on LAEs beyond $z>7$ is far less clear. Apart from a few spectroscopically confirmed LAEs at $z\\sim7$ (one spectroscopically confirmed out of two at $z=6.96$ \\citep{OTA08} and one spectroscopically confirmed out of three at $z=7.22$ \\citep{SHIBUYA12}), there are no confirmed LAEs at higher redshifts. A total sample of $\\sim15$ candidate LAEs at $z=7.7$ is known \\citep{HIBON10,TILVI10,KRUG12}. \\citet{TILVI10} and \\citet{KRUG12} favor a non-evolution of the Ly$\\alpha$ LF between $5.7 < z < 7.7$ (see also \\citet{HIBON11}), which is in tension with other narrow band searches for LAEs at $z > 7$ that only place limits on the number counts of LAEs \\citep{SOBRAL09,CLEMENT12,OTA12,MATTHEE14}. The reason for this tensions may be low-redshift interlopers and false detections in the LAE samples. At $z<7$, both of these are estimated to contribute less than 10--20\\% (see e.g., \\citet{OUCHI10}), at higher redshifts, these contribution are not known, yet, but are probably much higher (see \\citet{MATTHEE14} and this work). Spectroscopic follow-up observations of high redshift candidate LAEs are therefore necessary to resolve the tensions between the LAE and LBG results at $z>7$ and to constrain the process of re-ionization at higher redshifts.   In this work\\footnote{Magnitudes are given in the AB system and we assume a flat universe with $\\Omega_{m}=0.25$, $\\Omega_{\\Lambda}=0.75$, and H$_{0}=70$ km s$^{-1}$ Mpc$^{-1}$.}, we present Keck-I MOSFIRE spectroscopic follow-up of two $z\\sim7.7$ LAE candidates ($\\S2$) originally found by \\citet{KRUG12}. We then go on to compare these results to existing data at lower redshift ($\\S3$) and to an empirical model derived from the LBG UV continuum LF and observed equivalent width distribution ($\\S\\S4.1,4.2$). This allows us to place limits on the neutral fraction of the IGM at $z\\sim8$ ($\\S4.3$) and enables us to predict the LAE LF at $z\\sim8-9$ ($\\S5$).    \\begin{figure*} \\includegraphics[scale=.25,angle=0]{f1.pdf} \\caption{SUBARU z$^{+}$-band images centered on LAE1 (top) and LAE2 (bottom) overlaid with the MOSFIRE slits configuration (left). Observed (center) and simulated 2D-spectra (right) are shown as well, both are binned to obtain R $\\sim1500$. The wavelength range where the emission is expected from the UNB observations is marked with red lines. For the simulation shown here, we assumed a rest-frame FWHM of 1.5$\\textup{\\AA}$ for the Ly$\\alpha$ line (represented as truncated gaussian) and a spatial extent of 1$\\arcsec$. This simulation shows the clear detection of the line for both LAEs. \\label{fig:Slit}} \\end{figure*} ",
        "conclusions": "We have presented follow-up observations on two bright LAE candidates at $z\\sim7.7$ using MOSFIRE. We rule out any line emission at a level of several $\\sigma$ for both objects. The limits inferred from these non-detections suggest a strong evolution of the LAE LF between $6 < z < 8$, consistent with what is seen in LBG samples. We create an empirical model using the observed LBG UV continuum LFs and Ly$\\alpha$ rest-frame equivalent width distributions to understand the interplay between LAE and UV continuum selected galaxies. We find that our model and the observed LAE LF follow each other, but note a secondary effect which is due to a change in the EW$_{0}$ distribution of the galaxies as a function of redshift. From this differential evolution and assuming two different models on Ly$\\alpha$ absorption, we find consistent lower limits on the neutral hydrogen fraction at $z\\sim7.7$ of 50-70\\%. Furthermore, we find a strong evolution in the Ly$\\alpha$ optical depth at $z>6$ which can be characterized by $(1+z)^{2.2\\pm0.5}$. All in all, our results are indicative of a continuation of strong evolution in the IGM beyond $z=7$."
    },
    "1402/1402.1437_arXiv.txt": {
        "abstract": "{ Primordial black holes (PBHs) are of special interest because of the crucial role of quantum effects in their formation and evaporation. This means that they provide a unique probe of the early universe, high-energy physics and quantum gravity.  We highlight some recent developments in the subject, including improved limits on the fraction of the Universe going into evaporating PBHs in the mass range $ 10^9-10^{17}\\,\\mathrm g $ and the possibility of using PBHs to probe a cosmological bounce. } ",
        "introduction": "\\label{sec:1}  A comparison of the cosmological density at a time $t$ after the big bang with the density associated with a black hole of mass $M$ shows that PBHs should have of order the particle horizon mass, $M_H(t) \\approx 10^{15}(t/10^{-23}s)$g, at formation. They could thus span an enormous mass range: from $10^{-5}$g for those formed at $10^{-43}$s to $10^5M_{\\odot}$ for those formed  at 1~s. By contrast, black holes forming at the present epoch could never be smaller than about $1M_{\\odot}$. However, the high density of the early Universe is  not a {\\it sufficient} condition for PBH formation. One either needs large-amplitude density fluctuations, possibly of inflationary origin, so that overdense regions can eventually stop expanding and recollapse, or some sort of cosmological phase transition at which PBHs can form spontaneously (eg. via  the collapse of cosmic loops or the collisions of  bubbles of broken symmetry). All these formation mechanisms depend in some sense  on quantum effects and they are discussed in detail in Ref.~\\cite{cksy} (henceforth CKSY).  The realization that PBHs might be small prompted Hawking to study their quantum consequences This led to his famous discovery \\cite{h74} that black holes radiate thermally with a temperature $T \\approx 10^{-7}(M/M_{\\odot})^{-1}$K and evaporate on a timescale $ \\tau(M) \\approx 10^{64}(M/M_{\\odot})^3$y. Only black holes smaller than $10^{15}$g would have evaporated by the present epoch and $10^{15}$g ones would be exploding today. Since the latter would be producing photons with energy of order 100~MeV, the observational limits on the $\\gamma$-ray background intensity imply that their density could not exceed $10^{-8}$ times the critical density. Nevertheless, this does not preclude PBHs playing other important cosmological roles. Indeed, their study provides a unique probe of four areas of physics: gravitational collapse for $M>10^{15}g$, high energy physics for $M\\sim 10^{15}g$, the early Universe for $M<10^{15}g$ and quantum gravity for $M\\sim10^{-5}g$.  Since both their formation and evaporation are a consequence of quantum effects, PBHs may offer the only astrophysical realization of what might be termed ``quantum black holes'' (i.e. holes for which quantum effects are important)  \\cite{calmet}. This article will focus on their evaporation rather than their formation. In particular, it will discuss the upper limit on the fraction of the Universe going into PBHs as a function of mass  because this provides important constraints on models (such as inflation) predicting their formation. The fraction of the Universe collapsing into PBHs at time $t$ is related to their current density parameter $\\Omega_{\\rm PBH}$ by \\begin{equation} \\beta \\approx 10^{-6} \\Omega_{\\rm PBH} (t/s)^{1/2} \\approx 10^{-18}\\Omega_{\\rm PBH}(M/10^{15}g)^{1/2} \\end{equation} where the $t$ dependence reflects the decreasing ratio of the PBH and radiation densities at early times \\cite{c75}. Any limit on $\\Omega_{\\rm PBH}$ therefore places a constraint on $\\beta$ as a function of $M$. The constraints on $ \\beta(M) $ have been studied by numerous authors but the most recent and comprehensive discussion is that of Ref.~\\cite{cksy}. The limits cover the mass range $ 10^9-10^{17}\\,\\mathrm g $ and are shown in Fig.~\\ref{fig:combined}. The important point is that the value of $ \\beta(M) $ must be tiny throughout this mass range, so any cosmological model which entails an appreciable fraction of the Universe going into PBHs is immediately excluded. The most stringent limits -- associated with big bang nucleosynthesis (BBN), the extragalactic $\\gamma$-ray background (EGB) and observations of anisotropies  in the cosmic microwave background (CMB) -- are discussed below. Positive evidence for PBHs might come from cosmic rays or short-period gamma-ray bursts but this is not covered below since the status of the observations is still ambiguous. \\begin{figure}[b] \\sidecaption \\includegraphics[scale=.59]{combined.pdf} \\caption{ Combined BBN and EGB limits (solid), compared to other constraints on evaporating PBHs from LSP relics and CMB distortions (short-dashed), extragalactic antiprotons and neutrinos (dotted), the Galactic $ \\gamma $-ray background (long-dashed), CMB anisotropies (dash-dotted) and the density limit from the smallest unevaporated black holes (dashed). From Ref.~\\cite{cksy}. \\label{fig:combined} } \\end{figure} ",
        "conclusions": ""
    },
    "1402/1402.4815_arXiv.txt": {
        "abstract": "We present the observational evidence of the 2-halo term in the stacked shear profile of a sample of $\\sim 1200$ optically selected galaxy clusters based on imaging data and the public shear catalog from the CFHTLenS. We find that the halo bias, a measure of the correlated distribution of matter around galaxy clusters, has  amplitude and  correlation with galaxy cluster mass in very good agreement with the predictions based on the LCDM standard cosmological model. The mass-concentration relation is flat but higher than theoretical predictions. We also confirm the close scaling relation between the optical richness of galaxy clusters and their mass. ",
        "introduction": "A fundamental scope of the observational study of galaxy clusters, the most massive gravitationally bound systems in the cosmos, is testing models of structure formation and gaining insight into the complex physical interplay between the baryonic and dark components of matter. Currently, the most successful framework is based on the cosmological constant term and the cold, collisionless dark matter (LCDM), which nature is still not understood. Numerical simulations have provided clear predictions for the structure of dark matter halos (e.g., the so-called mass-concentration relation, \\citealt{Bullock2001}) and their clustering properties.  The halo power spectrum $P_{\\rm h} (k)$ describes the statistical properties of the spatial distribution of massive halos. Observed structures are in principle differently distributed from the underlying dark matter. The correlation between the halo density field and the underlying matter distribution is parametrized by the halo bias parameter, $\\bh$. This is defined as  the ratio of the halo power spectrum to the linear matter \\citep{Tinker2010, OguriTakada2011}: \\begin{equation} b^2_{\\rm h} (k)= \\frac{P_{\\rm h}(k)}{P_{\\rm m}(k)} \\, . \\label{eq:bias_def} \\end{equation} Theoretical models of gravitational clustering of dark matter haloes % provide robust predictions for the bias \\citep{Mo1996, Sheth2001}. It is expected to be constant at large scales and an increasing function of the  peak height. It is therefore of importance to compare such predictions with observations of galaxy clusters, as a function of mass and cosmic epoch.  The correlated matter distribution around galaxy clusters manifests itself as a 2-halo term in the mass density profiles, at scales lager than $\\sim 10 \\, {\\rm Mpc}$, its contribution being proportional to $\\bh$. The 2-halo term describes the cumulative effects of the large scale structure in which galaxy clusters are located. While weak gravitational lensing is a powerful tool to measure mass distribution in galaxy clusters \\citep{KneibNatarajan2011}, the lensing signal expected in the 2-halo term is too low to be reliably detected in the outskirts of individual systems. Stacking the lensing signal of a large number of galaxy clusters enable us to detect at high S/N the 2-halo term. Stacked shear profiles of galaxy groups and  clusters have been used to study the average properties of radial profiles and the observable scaling relations \\citep{Sheldon2001, Dahle2003, Okabe2010, Leauthaud2010, Oguri2012MNRAS, Sereno2013, Shan2013}.  \\cite{Johnston2007} used  $\\sim~1.3\\times 10^5$ galaxy groups and clusters from the Sloan Digital Sky Survey (SDSS) to obtain stacked lensing signals and infer the relations between mass and optical richness and luminosity over a wide range of masses, and  also to measure the halo bias and  concentration as a function of  mass. \\cite{OguriTakada2011} have proposed a cosmological test based on the combination of stacked weak lensing of galaxy clusters with number counts and correlation functions, showing that this reduces systematic errors from poorly known redshift distributions and scaling relations.  Several studies have shown observational evidence for the 2-halo term in the stacked galaxy-galaxy lensing profile \\citep{Brimioulle2013, Velander2013}. In this work, we apply this technique at the galaxy clusters scale. We use public data from CFHTLenS to measure the lensing signal from a sample of optically selected galaxy clusters \\citep{Wen2012}. As framework we use a spatially flat, cosmological model with $\\Omega_\\Lambda  = 0.7$, $\\Omega_{\\rm m} = 0.3$ and $h=0.7$, with $H_0 = 100 \\, h \\, \\kms \\, {\\rm Mpc}^{-1}$. ",
        "conclusions": "We  used the CFHTLenS public shear catalog to obtain the stacked weak lensing signal of 1176 galaxy clusters in the redshift range $0.1 < z < 0.6 \\, $. With respect to previous observational studies of the galaxy cluster halo bias based on SDSS imaging data, we used shear measurements derived from the deeper and higher quality data from the CFHTLS. We have built high S/N profiles in six bins in optical richness $R_{L_*}$, which average masses span the range from 0.5 to $\\sim 3 \\, \\times 10^{14} \\, M_{\\odot}/h$.  We confirm that the cluster mass robustly correlates  with the optical richness, with parameters in agreement with those found by \\cite{Wen2012}. Our result supports the use of the stacking method in weak lensing to calibrate galaxy clusters scaling relations, e.g. \\cite{OguriTakada2011} .  We studied the $c(M)$ and the $\\bh (M)$ relations, and we found in both cases an agreement with theoretical predictions based on the LCDM cosmological models. While consistent at 1-$\\sigma$ level with the predictions from \\cite{Duffy2008}, we find evidence for an over-concentration of the $c(M)$-relation, in closer agreement with recent simulations by \\cite{Bhattacharya2013}.  Our results partially reconcile present tension between observed $c(M)$ relations and theoretical predictions. On one hand, studies of single clusters found very steep and over-concentrated relations \\citep{Oguri2012MNRAS, fed12, Sereno2013}. On the other hand, previous stacked analyses found flat relations of the expected amplitude (Mandelbaum et al. 2008). We found an alternative scenario on a middling ground: an over-concentrated but still flat relation. Analysis of individual clusters might be affected by low S/N and high correlation, as well as by selection effects whereas previous SDSS stacked analysis might suffer from contamination and off-set effects.  The large scale shear profile is a degenerate function of the halo bias and the power spectrum normalisation, therefore our measurement of the halo bias term is not independent of  $\\sigma_8$. Our measurement of the quantity $\\bh (M)\\, \\sigma_8^2$ as a function of the average halo mass is in very good agreement with theoretical predictions by Tinker et al. (2010), and systematically higher than the ones by \\cite{Johnston2007}, Fig.~\\ref{fig:bias}. As shown by \\cite{Rozo2010}, shear measurements in \\cite{Johnston2007} are underestimated by $\\sim$ 20\\% due to the diluted lensing signal by foreground galaxies. This systematic error might explain the discrepancy with our measurements.  Our analysis is also consistent with measurements from the correlation function of galaxy clusters. \\citet{Veropalumbo2013} analyzed a spectroscopic sample of 25226 clusters of richness $R_{L_*}\\geq 12$ at $z\\sim 0.3$ and obtained $\\bh=2.06\\pm0.04$ for $\\sigma_8=0.8$.  Our results open the possibility to verify further physical effects on the clustering properties of massive halos: \\cite{Villaescusa2013} used $N$-body simulations to show that cosmological models with massive neutrinos show a scale-dependent bias on large scales, while  \\cite{Dalal2008} have shown that the non-Gaussianity of primordial fluctuations brings to a strongly scale-dependent bias."
    },
    "1402/1402.0006_arXiv.txt": {
        "abstract": "The total infrared (IR) luminosity is very useful for estimating the star formation rate (SFR) of galaxies, but converting the IR luminosity into an SFR relies on assumptions that do not hold for all galaxies. We test the effectiveness of the IR luminosity as an SFR indicator by applying it to synthetic spectral energy distributions generated from three-dimensional hydrodynamical simulations of isolated disc galaxies and galaxy mergers. In general, the SFR inferred from the IR luminosity agrees well with the true instantaneous SFR of the simulated galaxies. However, for the major mergers in which a strong starburst is induced, the SFR inferred from the IR luminosity can overestimate the instantaneous SFR during the post-starburst phase by greater than two orders of magnitude. Even though the instantaneous SFR decreases rapidly after the starburst, the stars that were formed in the starburst can remain dust-obscured and thus produce significant IR luminosity. Consequently, use of the IR luminosity as an SFR indicator may cause one to conclude that post-starburst galaxies are still star-forming, whereas in reality, star formation was recently quenched. ",
        "introduction": "\\label{S:intro}  The total infrared (IR) luminosity ($\\lir$) of galaxies has been used to infer the star formation rate (SFR) of galaxies for decades (see \\citealt{Kennicutt:1998review} and \\citealt{Kennicutt:2012} for reviews). Since the launch of the \\textit{Herschel Space Observatory}, the IR luminosity has been often used to infer the SFRs of diverse types of local and high-redshift galaxies \\citep[e.g.][]{Rodighiero:2010,Elbaz:2011,Wuyts:2011a, Magnelli:2012,Rosario:2012,Rosario:2013,Lee:2013}.  Using the IR luminosity as an SFR indicator has several advantages. For example, other SFR indicators, such as recombination lines (e.g. \\halpha) and ultraviolet (UV) emission, must be corrected for dust attenuation to recover the intrinsic SFR. Correcting for dust is inherently uncertain because the observed emission is dominated by the less-obscured lines of sight; consequently, the dust attenuation can be underestimated. In contrast, the IR is a promising SFR indicator precisely because of dust: in the limit of high obscuration, essentially all of the UV-optical light from young stars is absorbed and re-emitted into the IR. Thus, the IR luminosity should be the best SFR indicator in those situations, such as for starburst galaxies. Furthermore, the IR luminosity can be combined with tracers of unobscured star formation (uncorrected for dust attenuation); such combinations may yield more accurate SFR estimates (e.g. \\citealt{Kennicutt:2007,Kennicutt:2009}; \\citealt*{Relano:2009}; \\citealt{Wuyts:2011a,Wuyts:2011b,Reddy:2012, Lanz:2013}).  Naturally, there are also caveats to using the IR luminosity as an SFR indicator because multiple assumptions must be made to convert the IR luminosity into an SFR (see \\citealt{Kennicutt:1998review} for a more thorough discussion). As with all SFR indicators, the IR luminosity is sensitive to the initial mass function (IMF; e.g. \\citealt{Narayanan:2012IMF}). The sensitivity to the high end of the IMF is less severe for the IR than for recombination lines because the IR light originates from dust that is heated by both lower-mass stars and the massive stars that photoionise nebulae and thereby generate recombination line emission. However, this effect can also be a disadvantage of using the IR luminosity as an SFR indicator: it is commonly stated \\citep[e.g.][]{Kennicutt:1998review} that, similar to the UV, the IR luminosity is sensitive to the SFR averaged over the past $\\sim 100$ Myr (in contrast, recombination lines are sensitive to the SFR averaged over the past $\\sim 10$ Myr). Thus, the IR luminosity may not trace the instantaneous SFR accurately if it varies over timescales of $\\la 100$ Myr.  Furthermore, dust can also be heated by older stars; if such dust heating is non-negligible, as has been suggested by some authors (e.g. \\citealt*{Sauvage:1992}; \\citealt{Smith:1994}; \\citealt*{Smith:1996,Walterbos:1996}; \\citealt{Kennicutt:2009,Salim:2009}; \\citealt{Calzetti:2010}; \\citealt{Kelson:2010,Murphy:2011,Totani:2011}; \\citealt{Groves:2012, Leroy:2012}; \\citealt{Fumagalli:2013,Utomo2014}) converting the IR luminosity into an SFR using a standard calibration will overestimate the true SFR. This effect becomes increasingly important with decreasing SFR \\citep{Calzetti:2010,Smith:2012, Rowlands2014,Utomo2014}. When the contribution of older stellar populations to dust heating is non-negligible, the implicit averaging timescale is $>> 100$ Myr \\citep{Salim:2009,Kelson:2010}. Emission from active galactic nuclei (AGN) can also {heat dust and thus cause the IR-inferred SFR to overestimate the true SFR. The IR-inferred SFR can also under-predict the true SFR if a significant fraction of the photons from young stars is not absorbed by dust or if stellar populations younger than 100 Myr dominate the bolometric luminosity \\citep[e.g.][]{Schaerer:2013, Sklias2014}.  In this paper, we quantify the effectiveness of the IR luminosity as an SFR indicator by applying it to mock spectral energy distributions (SEDs) calculated from three-dimensional hydrodynamical simulations of isolated disc galaxies and galaxy mergers. Because the simulated galaxies have star formation histories that are significantly more complex than those assumed to calibrate the $\\lir$-SFR relation (i.e., a single burst or continuous star formation history), realistically treat dust attenuation from both stellar birth clouds and the diffuse interstellar medium (ISM), and include AGN emission, they can be used to characterise the effects of some of the aforementioned assumptions that may cause the IR luminosity to poorly trace the instantaneous SFR. The remainder of this work is organised as follows: in Section \\ref{S:methods}, we briefly summarise the details of the simulations. In Section \\ref{S:results}, we compare the $\\lir$-inferred SFR and true instantaneous SFR and the resulting inferred and true SFR--stellar mass ($\\mstar$) relations. In Section \\ref{S:discussion}, we discuss some implications of our results and ways forward. ",
        "conclusions": "We have used a set of mock SEDs generated by performing dust radiative transfer on hydrodynamical simulations of isolated disc galaxies and galaxy mergers to investigate the effectiveness of the IR luminosity at recovering the true instantaneous SFR of the simulated galaxies. Our principal conclusions are the following: \\begin{enumerate} \\item For most galaxies (i.e., quiescently star-forming disc galaxies and minor mergers), the SFR inferred from $\\lir$ agrees very well with the true SFR. \\item In the simulations in which a strong starburst occurs, $\\lir$ can significantly overestimate the SFR after the peak of the starburst. The reason for this overestimation is that although the SFR decreases rapidly from the peak, there is still significant IR emission from dust heated by the stars formed during, and even before, the starburst. \\item The magnitude of the overestimate is greater for lower SSFR values. However, the instantaneous SSFR value is not the sole determinant of the contribution of older stellar populations to the dust heating because for fixed SSFR, the overestimation is more severe in the post-starburst phase of the simulations. Thus, the time evolution of the SSFR, not just the current value, is an important determinant of how well the IR-inferred SFR traces the true SFR. \\item This overestimation may have significant implications for e.g. the SFR-stellar mass relation. In particular, it may cause the number of quenched galaxies and degree to which galaxies are quenched to be underestimated. \\end{enumerate} The results presented in this work highlight the need for caution when applying simple SFR tracers to galaxies for which the underlying assumptions may not hold. More-sophisticated techniques, such as SED modelling, should be used when possible, but one must also mind the caveats and potential inherent degeneracies in these techniques."
    },
    "1402/1402.0695_arXiv.txt": {
        "abstract": "{On the basis of a fermionic dark matter model we fit rotation curves of The HI Nearby Galaxy Survey THINGS sample and compare our 3-parametric model to other models widely used in the literature: 2-parametric Navarro--Frenk--White, pseudoisothermal sphere, Burkhert models, and 3-parametric Einasto model, suggested as the new \"standard dark matter profile\" model in the paper by Chemin et. al., AJ 142 (2011) 109. The results from the fitting procedure provides evidence for an underlying fermionic nature of the dark matter candidate, with rest mass above the keV regime.} ",
        "introduction": "The problem of the distribution of Dark Matter (DM) in galaxies, as usually addressed in the literature, is mainly focused in the halo regions and associated with the galaxy rotation curves obtained from the observations, see e.g. \\cite{2013BrJPh..43..369E}. A well-known approach used to deal with this problem is the Navarro--Frenk--White (NFW) model \\citep{1997ApJ...490..493N}, expected to provide a universal description of dark matter halos obtained under the following main considerations: $1)$ N-body simulations in Cold dark matter (CDM) and ($\\Lambda$CDM) cosmologies; $2)$ particles each of masses of $\\sim 10^9 M_{\\odot}$ \\footnote{Modern numerical simulations can reach better resolution down to particle masses of $\\sim 10^5 M_{\\odot}$ \\citep{2012MNRAS.425.2169G}.}; $3)$ classical Newtonian physics.  Despite an indicated agreement of this model with the large scale structure of the Universe, some problems remains at galactic scales, see e.g. \\citet{2013ApJ...766...56M}. A central characteristic of the NFW dark matter density profiles, is that they show a cuspy and divergent behaviour through the center of the configuration, while empirical profiles tend to show a core of constant density, giving rise to the well-known core-cusp controversy, see e.g. \\citet{2010AdAst2010E...5D}.  Yet another important approach developed to understand the distribution of matter in galaxies has been advanced by \\citet{1965TrAlm...5...87E} and \\citet{1989A&A...223...89E}. This is a phenomenological approach consisting in the proposal of an empirical fitting function composed by three free parameters as detailed in equation (\\ref{eq:2}) \\begin{equation} \\rho_E(r)=\\rho_{-2}\\exp{\\left(-\\frac{2}{n}\\left[\\left(\\frac{r}{r_{-2}}\\right)^n-1\\right]\\right)}, \\label{eq:2} \\end{equation} where $\\rho_{-2}$ and $r_{-2}$ are the density and radius at which $\\rho(r)\\propto r^{-2}$, and $n$ is the Einasto index which determines the shape of the profile.  Recent N-body simulations in $\\Lambda$CDM cosmology by \\citet{2004MNRAS.349.1039N} purported a novel dark matter halo model different from NFW. This model was soon realized \\citep{2006AJ....132.2685M} to be the same as the Einasto one as given by equation (\\ref{eq:2}).  After that, using the highest quality rotation curves available to date obtained from The HI Nearby Galaxy Survey (THINGS) \\citep{2008AJ....136.2563W, 2008AJ....136.2648D}, the Einasto dark matter halo model has been proposed as the standard model for dark matter halos by \\citet{2011AJ....142..109C}, as it provides both cored and cusped distributions for different values of model parameters (see Fig.~\\ref{fig1}). In that work, the fundamental core-cusp discrepancy is analyzed in detail for the whole sample of galaxies under study. It is clearly shown that for the majority of the galaxies considered in the sample, the cored halos (compatible with near unity Einasto indexes) are preferred over the cuspy ones (these instead compatible with higher Einasto indexes).  \\begin{figure}[!h] \\centerline{\\includegraphics[width=0.85\\columnwidth]{Fig11}} \\vspace*{8pt} \\caption{Comparison of the density profiles for different phenomenological models of dark matter distribution.}\\label{fig1} \\end{figure}  We present here a novel approach focusing on galactic structures and an underlying microphysical component of Dark Matter. The model is built upon the following general considerations: $1)$ the Dark Matter component is assumed to be chargeless spin-1/2 fermions; $2)$ the configurations are described by General Relativity; $3)$ the particles are assumed to be isothermal in thermodynamic equilibrium (i.e. without the need of pre-fixing any cosmological history). The theoretical fundament of this new approach is detailed in the model of semi-degenerate self-gravitating fermions first introduced by \\citet{1990A&A...235....1G}, and more recently with applications to galactic dark matter by \\citet{2013PRL..R, 2013JKPS..A}. Our model is based in the following main assumptions: \\begin{enumerate} \\item the problem of galactic cores and halos have to be addressed unitarily; \\item for definiteness we study the simplest problem of \"bare\" massive particles, neglecting at this stage all other interactions than the gravitational one and fulfilling only the Fermi--Dirac statistical distribution \\begin{equation}% f=\\frac{1}{\\exp\\left({\\frac{\\epsilon-\\mu}{kT}}\\right)+1}= \\frac{1}{\\exp{\\left(\\frac{\\epsilon}{\\beta mc^2}-\\theta \\right)}+1}, \\end{equation} where $\\epsilon$ is kinetic energy of the particles, $\\mu$ is chemical potential, $T$ is the temperature, $k$ is Boltzmann constant and $c$ is the speed of light. The mass of the particle ($m$), the temperature parameter ($\\beta=kT/mc^2$) and the degeneracy parameter ($\\theta=\\mu/kT$) at the center are the three free parameters of the model; \\item we consider zero total angular momentum and also neglect any effect of baryonic matter on the DM in the mathematical formulation. \\end{enumerate}  It is shown that in any such system the density at large radii scales as $r^{-2}$ independently of the values of the central density, providing the flat rotation curve \\citep{1990A&A...235....1G, 2013PRL..R, 2013JKPS..A}.  The dark matter halos obtained in the new dark matter approach proposed here share a common or universal feature which shed more light on the core-cusp discrepancy, while providing a new mass scale to the dark matter candidate. Our density profiles always favor a cored behaviour (without any cusp) in the observed inner halo regions, given the quantum nature of the fermionic particles. Another fundamental outcome of our model is the range of the DM particle mass, which must be $m\\gtrsim5$ keV in order to be in agreement with typical halo sizes of the observed dwarf galaxies \\citep{2013PRL..R}.  The theoretical formulation of \\citet{2013PRL..R, 2013JKPS..A} is based on the first principles physics and provides a physical complement to Einasto phenomenological models. It also offers the necessity to approach the Dark Matter distribution in galactic haloes with fermions with masses larger than the above mentioned bound.  \\begin{figure}[!h] \\centerline{\\includegraphics[width=0.85\\columnwidth]{Fig12}} \\vspace*{8pt} \\caption{Semidegenerate density profile in dimensionless units for central degeneracy parameter $\\theta_0=15$ and central temperature parameter $\\beta_0=10^{-10}$.}\\label{fig12} \\end{figure}  The paper is structured as follows. We model the distribution of Dark Matter as semidegenerate fully relaxed thermal self-gravitating general relativistic fermionic solutions of \\citet{1990A&A...235....1G}, see Sec.~\\ref{model}. The resulting density profiles provide flat rotation curve at large distances, cored distribution of dark matter in the halo, and a massive degenerate core at the very center, see Sec.~\\ref{solutions}. We describe in Sec.~\\ref{procedure} the actual procedure of fitting of rotation curves, and then discuss the results in Sec.~\\ref{results}. Conclusions follow. ",
        "conclusions": "It follows from the results of fitting that the semidegenerate fermionic distributions can fit dark matter in the THINGS sample of galaxies at least as well as other profiles considered in the literature, with the important \"revenue\" that this profile is theoretically motivated, and is not phenomenological as most of the others. The cases when Einasto profile fits rotational curve much better than semidegenerate profile show the general cuspy behaviour of dark matter distribution, possibly representing a special class of galaxies that are still not completely relaxed.  While Einasto profile is a pure phenomenological one based on best fit of the observational and numerical simulation data, our profile is derived from the first principles and based on the General Relativistic treatment of self-gravitating neutral fermions. There is the distinct possibility that our treatment gives the conceptual physical motivation for the existence of the cored Einasto profile directly from the structure of the microphysical constituents of dark matter."
    },
    "1402/1402.0376_arXiv.txt": {
        "abstract": "{ We investigate the contribution that a local over- or under-density can have on linear cosmic dipole estimations. We focus here on radio surveys, such as the NRAO VLA Sky Survey (NVSS), and forthcoming surveys such as those with the LOw Frequency ARray (LOFAR), the Australian Square Kilometre Array Pathfinder (ASKAP) and the Square Kilometre Array (SKA). The NVSS has already been used to estimate the cosmic radio dipole; it was shown recently that this radio dipole amplitude is larger than expected from a purely kinematic effect, assuming the velocity inferred from the dipole of the cosmic microwave background. We show here that a significant contribution to this excess could come from a local void or similar structure. In contrast to the kinetic contribution to the radio dipole, the structure dipole depends on the flux threshold of the survey and the wave band, which opens an opportunity to distinguish the two contributions.} ",
        "introduction": "In recent years,  the dipole anisotropy in radio surveys, such as the NVSS catalogue \\citep{NVSS}, has been investigated (e.g. \\citet{BW02}, \\citet{Singal11}, \\citet{Gibelyou}, \\citet{firstpaper} and \\citet{Jain13}). It appears that the cosmic radio dipole has a similar direction to the one found in the Cosmic Microwave Background (CMB), but with a significantly higher amplitude (by a factor of three to four, based on different estimators and surveys \\citep{Singal11, firstpaper}). In this work we investigate one possible effect which can increase the dipole amplitude observed in radio surveys, with respect to the CMB dipole.  There have recently been studies (e.g. \\citet{Void1} and \\citet{Void2}) which claim that the local universe (i.e. on scales of 300 Mpc) has an atypically low density of galaxies. If we do live in such a region, what would we expect to see regarding the observed cosmic radio dipole? We are unlikely to be living in the very centre of such a void, so  there will be some offset distance between us and the centre of the void, which we call $\\vec r_{\\rm v}$. If we imagine a sphere around the observer (in our case the Local Group), with a radius $R_{\\rm o}$ greater than the void radius $R_{\\rm v}$, we will expect to see more galaxies in one direction than in the other.  It is likely that the Local Group moves towards the direction where we see more galaxies, due to their gravitational pull. This direction has been determined to be $(l,b)$ = ($276^\\circ \\pm 3^\\circ, 30^\\circ \\pm 3^\\circ$) \\citep{Kogut} in galactic coordinates. The CMB dipole, $(l,b)$ = ($263.99^\\circ \\pm 0.14^\\circ, 48.26^\\circ \\pm 0.03^\\circ$) from \\citet{CMB_dipole_obs}, is caused by the motion of the Sun relative to the CMB, while the radio dipole, $(l,b)$ = ($248^\\circ \\pm 28^\\circ, 46^\\circ \\pm 19^\\circ$) from \\citet{firstpaper}, can be expected to receive contributions from the motion of the Solar System with respect to the CMB (kinetic dipole) and due to the uneven galaxy distribution (structure dipole). Within the current  accuracy, the direction of the radio dipole agrees with the CMB direction as well as with the motion of the Local Group with respect to the CMB. Therefore we expect the contribution of a local void to the radio dipole to add up with the velocity dipole, resulting in a larger dipole amplitude in radio surveys.  The local structures considered in this work are not in conflict with the Copernican principle, as they are much smaller than the Hubble scale and thus a fine tuning of the position of the observer with respect to the centre of a void is not required. This is different to scenarios in which huge voids have been invoked to provide an alternative explanation of dark energy (e.g. \\citet{Celerier}, \\citet{Alnes05}, \\citet{Alnes06}).   In this work, we will investigate this chain of thought in a more quantitative manner. Our model will be discussed in section \\ref{Model}, followed by detailed testing in section \\ref{Testing}. In section \\ref{MissingDipole} we will examine the effects of realistic voids on the dipole, and we will present our conclusions in section \\ref{Conclusion}. ",
        "conclusions": "\\label{Conclusion}  We have been able to develop a model that can describe the influence of spherically symmetric local structures on linear dipole estimators. This model was tested and confirmed by computer simulations to a high level of accuracy. From this model we learn how the structure parameters (void size $R_{\\rm v}$, observer distance from the void centre $r_{\\rm v}$ and void density contrast $\\delta$) influence the structure dipole amplitude.  Our analytical model requires a constant background density, which is a reasonable approximation at small redshifts. In order to include the effects of cosmic expansion and galaxy evolution, the dipole contribution for the realistic void model was estimated by means of simulations of the number counts of radio galaxies. Not included in this work are estimations of the dipole contribution from several galaxy clusters or other structures; because the dipole estimator is linear, these would just be a sum of terms similar to the ones calculated in this paper.  One might ask if a void as considered in this work would show up in the CMB. It is clear that the size of the effect will depend on the distance of the observer from the centre of the void. The CMB dipole will be maximally affected for an observer sitting at the edge of the void.  There is no contribution to the CMB dipole if the observer sits in the centre and the effect would show up at much smaller angular scales if the observer were far away form the void. A CMB dipole would be induced by the integrated Sachs-Wolfe effect (or the Rees-Sciama effect if non-linear effects play a role) and thus the dynamics of the void profile would be important to determine its amplitude. Such structures have been studied previously, e.g. by \\citet{ISWVoid1}, \\citet{ISWVoid2}, \\citet{ISWVoid3}, \\citet{ISWVoid4}, \\citet{ISWVoid5}, \\citet{ISWVoid6} and \\citet{ISWVoid7}. An order of magnitude estimate of the maximum possible effect \\citep{ISWVoid8} gives \\begin{equation} \\Delta T/T \\sim \\delta^{3/2} \\left(\\frac{R_{\\rm v}}{R_{\\rm H}} \\right)^3 \\sim 2 \\times 10^{-4} \\end{equation} for the model considered here. This shows that the void-like structure considered in this paper would not be in tension with the observed CMB dipole, but might contribute to the CMB anomalies at small multipole moments. A detailed study of this topic is beyond the scope of this work.  For the void model of \\cite{Void1} for our local environment, we have run  simulations which include a radio sky model from \\cite{SSS}. We found  that such a void already has a significant effect on the dipole estimation for surveys like the NVSS. The dipole amplitude measured by the linear estimator from \\citet{firstpaper} of this void is expected to be $\\tilde d_{\\rm{void}} = 0.21 \\pm 0.01 \\ \\times 10^{-2}$. The discrepancy between radio and CMB dipole measurements can be relaxed by such a contribution, but the difference cannot be explained completely by the contribution from a single, realistic void. In forthcoming surveys, with lower flux density limits, the effect of local structure will become even more important."
    },
    "1402/1402.0971_arXiv.txt": {
        "abstract": "We present a high signal to noise (S/N$>$10)  medium resolution (R=2000) LBT-LUCIFER spectrum of the early-type galaxy (ETG) S2F1-142 at $z\\simeq1.4$. By means of the CaT line at $8662$ \\AA, we measured its redshift $z=1.386\\pm 0.001$ and we estimated its velocity dispersion $\\sigma_{v}=340  ^{-60}_{+120}$ km/s. Its corresponding virial mass is 3.9$\\times10^{11}$ M$_\\odot$, compatible with the stellar mass estimates obtained assuming Initial Mass Functions (IMFs) less dwarf rich than the Salpeter one. S2F1-142 is a compact galaxy with $R_{e}$=3.1$\\pm$0.2 kpc, i.e., an effective radius more than three times smaller than the average $R_{e}$ of early-type galaxies with the same mass in the local universe. At the same time, we found local and high redshift galaxies with a similar mass content and  similar effective radius confirming that it is fully consistent with the already available measures of $R_{e}$ and $\\sigma_{v}$ both in the local and in the distant universe. Considering the distribution of $R_{e}$ and $\\sigma_{v}$ as a function of the stellar mass content of ETGs, both in the local and in the distant universe,  we noticed that the measured velocity dispersions of the more compact galaxies are on average slightly lower than expected on the basis of their compactness and the virial theorem, suggesting that {\\it i)} their dark matter content is lower than in the more diffuse galaxies and/or {\\it ii)} their luminosity profiles are steeper than in the more diffuse galaxies and/or {\\it iii)} their larger compactness is an apparent effect caused by the overestimate of their stellar mass content (due to bottom lighter IMF and/or systematic affecting the stellar mass estimates). ",
        "introduction": "Understanding the formation of early-type galaxies (ETGs) is one of the crucial issues of cosmology, since they contain most of the present-day stars and baryons (e.g. \\citealt{ren}, \\citealt{fuk}). After their first spectroscopic detection at $z>1.5-2$ (\\citealt{dun}, \\citealt{spin}, \\citealt{s3}, \\citealt{mc}, \\citealt{cim04}, \\citealt{gla}), only in the last 10 years many estimates of their physical properties have been possible, e.g., age and metallicity of their stellar content (\\citealt{gobat}, \\citealt{l5}, \\citealt{daddi}), morphological profiles  and color gradients (\\citealt{gar12}, \\citealt{gar11}, \\citealt{mcg}, \\citealt{moth}). The emerging picture seems to tell us that most of them formed the bulk of their stars at $z>2-3$, even if further small star forming events happened during their successive evolution (e.g., \\citealt{ren}). Furthermore,  observations of $z\\simeq 1.5$ massive ETGs show evidence of both compact (i.e., $R_{e}$ a factor 3-6 smaller than local massive galaxies) and more diffuse galaxies (\\citealt{s10}, \\citealt{man10}), and the evolution  of the most compact ETGs from $z>1.5$ to $z=0$ in the sense of an enlarging of their effective radii is still a debated possibility (\\citealt{s9}; \\citealt{ono}; \\citealt{man10}). Within this picture, a still missing measure of high redshift ETGs properties is their velocity dispersion $\\sigma_{v}$, up to now available only for few $z>1.2$ galaxies and even for fewer compact galaxies (i.e., \\citealt{vdk}; \\citealt{toft}; \\citealt{vds}). The measure of the velocity dispersion allows to tackle the issue of ETGs formation and evolution from the point of view of their dynamics, and to estimate their total mass content, a primary parameter in galaxy formation models.     This paper presents a new velocity dispersion measure of a compact early-type galaxy (ETG) at $z\\sim$1.4, and it discusses the up to now available measures. The new $\\sigma_{v}$ measure is based on near-IR LBT-LUCIFER observations. Observations and data reduction are described in section 2. Section 3 presents details on the kinematic measures, which are finally discussed in section 4. We assume $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$=0.3 and $\\Omega_{\\Lambda}$=0.7. Magnitudes are in Vega system unless otherwise specified.       \\section[]{Observations and data reduction} The target is S2F1-142, a bright (K=17.6) ETG that has already been studied in details on the basis of its available photometric (MUNICS, \\citealt{d1}) and spectroscopic (\\citealt{l5}) data. Photometric available data included the V, R, I, J and K bands, while the low resolution spectroscopic data covered the range between 0.95 and 2.3 $\\mu$m. Furthermore, during HST CYCLE 14  we obtained deep HST-NICMOS high-resolution (0.075 arcsec/pix) imaging in the F160W band ($\\lambda\\sim1.6$ $\\mu$m) of this galaxy \\citep{l7}. The NICMOS high resolution near-IR data allowed us to measure its effective radius $R_{e}$ and light profile in the rest-frame R-band, revealing its compact nature (\\citealt{l7}, \\citealt{tru}). Table 1 summarizes the basic parameters and physical properties of  S2F1-142, while Figure 1 shows its HST F160W image  and its light profile.    \\begin{table*} \\centerline{ \\begin{tabular}{lccccccc} \\hline \\hline Object  & F160W$_{fit}^{(1)}$& M$_{R}$ & r$_e$ &  $R_{e}$ &$\\langle\\mu\\rangle_e^{F160W}$ & $\\langle\\mu\\rangle_e^R$ & K \\\\ &[mag]      &[mag]   & [arcsec] &  [kpc] & [mag/arcsec$^2$]& [mag/arcsec$^2$] & mag \\\\ \\hline S2F1-142 &18.65$\\pm$0.03  & -24.00 &  0.36$\\pm$0.02   & 3.1$\\pm$0.2  &   20.5$\\pm$0.2   &  18.4$\\pm$0.2  &17.8$\\pm$ 0.1   \\\\ \\hline \\hline \\end{tabular} } \\caption{Basic properties of S2F1-142. $(1)$: total integrated magnitude as extrapolated on the basis of the light profile fit.} \\end{table*}    \\begin{figure} \\includegraphics[width=4cm]{PaperSpec_fig1NEW.ps} \\includegraphics[width=4.5cm]{PaperSpec_fig1B.ps} \\caption{{\\it Left panel:} F160 image of S2F1-142, 3.7$^{''}$x3.7$^{''}$, total exposure time is 1h20m. {\\it Right panel:} observed (black points) and fitted (red and green lines) light profile of S2F1-142 in the F160 filter and residuals between observed and fitted profiles \\citep{l7}. Green and red lines report the results obtained assuming the Sersic and De Vaucouleurs profile, respectively.  } \\end{figure}    The spectroscopic observations of the target have been carried out with LUCIFER on the Large Binocular Telescope (LBT) at Mt. Graham (Arizona, USA) in long slit spectroscopic mode, during the semesters 2010A and 2010B.  \\noindent Observations during 2010A have been performed for 3 hours under good sky conditions (15th January and 18th February) and for other 45 minutes under poor transparency and bad seeing conditions (15th January). During data reduction we decided to discard the 45 minutes of observations made under bad sky conditions, since we verified that they do not contribute in increasing the final signal to noise.  \\noindent The N3.75 camera was used for this first set of observations, coupled with the 1$\\arcsec$  slit. The K$_{s}$ filter and the grating 150\\_K$_{s}$ have been adopted to cover the wavelength range 1.9$<\\lambda<$2.3 $\\mu$m with a sampling of 1.3\\AA\\ per pixel. The resulting spectral resolution is $R\\simeq$2000 corresponding to FWHM=10\\AA\\ ($\\sigma_{v}$=60 km/s).  \\noindent Observations carried out in 2010B added 9 hours of exposure on the target, all performed under good sky conditions. With respect to period 2010A, for this second set of observations we adopted the N1.8 camera that coupled with the same slit, grating and filter as before, results in the same spectral resolution but in a larger sampling (2.6\\AA) and a slightly larger wavelength coverage.  \\noindent All the observations have been performed by fixing the PA at -68.8 degree (North to East) in order to align the slit along the direction connecting our target S2F1-142 with a reference bright star (SDSS J030633.97-000242.8, K=15.2). This observational configuration has been chosen to help the slit positioning on our faint target. Furthermore, the bright reference star has been used during data reduction to correct for the final resulting sensitivity function (see below).  \\noindent Scientific observations have been split into many couples of short ($\\sim$ 5 minutes) exposures with the target located in two positions (A and B) along the slit (spaced $\\simeq 10$ arcsec).  \\noindent In each of the observing session we collected also some flat field images obtained with internal quartz lamp during daytime and used to build the final frame adopted to flatfield the scientific frames.   Images have been reduced by means of IRAF\\footnote{IRAF (Image Reduction and Analysis Facility) 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 following standard longslit reduction procedures. Wavelength calibration has been performed by means of the identification of the emission sky lines in each single frame. The row by row identification of the sky lines in each frame has also supplied the correction for their curvature along the spatial direction.  Wavelength calibration on each single frame results correct within 0.30\\AA\\ and 0.15\\AA\\ (rms with respect to the values of the identified lines) for the frames observed in 2010A and 2010B, respectively. A first step of background subtraction on each single frame has been performed by subtracting its associated frame of the dithering observing sequence A-B. The IRAF tool {\\it background} has been then used to refine the final sky subtraction of each frame, before aligning and coadding all of them. The final bi-dimensional spectral frame obtained in 9 hours of observations (i.e., 2010B run) is shown in Figure 2 (bottom). No telluric correction has been further applied, since after the previous background correction step no evident telluric absorption was appreciable.   Before extracting the one dimensional spectrum of the target, the final co-added images have been corrected for the response sensitivity functions. The latter have been derived by means of the bright reference star SDSS J030633.97-000242.8 (within the slit of all the scientific observations) compared with a standard spectrum of an M2 star taken from the Pickles Atlas \\citep{p8}. The SDSS J030633.97-000242.8 star has not been officially classified yet, but its colours in the optical and near-IR identify it as an M star. We choose to compare it with the M2 spectral template from the Pickles Atlas after verifying that results obtained by means of other M stars had negligible difference (well within the flux uncertainty due to the statistical noise) with the adopted one. Once reported to the same larger dispersion of 2.6\\AA\\ per pixel, the one dimensional spectra extracted from the two observing run have been summed to a single final spectrum. The one dimensional spectra have been extracted just adding the rows containing signal higher than 3 $\\sigma$ of noise, and following the spatial distortion along the slit.   Figure 3 shows the final resulting spectrum (gray and black lines) together with the residual associated noise (gray line in the lower part of the figure): the corresponding total S/N per pixel is $\\approx$ 12 in all the regions not affected by strong emission sky lines corresponding to about 7 \\AA\\ in the restframe.   \\begin{figure} \\includegraphics[width=9cm]{PaperSpec_fig2B.ps} \\includegraphics[width=9cm]{PaperSpec_fig2A.ps} \\caption{{\\it Top}: LBT-LUCIFER frame (5 minutes of exposure time). {\\it Bottom}: final reduced bi-dimensional spectrum (9 hours of exposure time). } \\end{figure}   \\begin{figure} \\includegraphics[width=8.8cm]{PaperSpec_fig3bis.ps} \\caption{Extracted spectrum of S2F1-142, corresponding to 12 hours of observations (light-gray line). The black (thick line) histogram shows the same spectrum binned over 4 pixels ($\\simeq 10$\\AA). The red line represents the bestfit template. Residual noise (i.e., the residual signal left after the sky subtraction procedure in a background region) is also reported (light gray line in the bottom part of the figure). The green line shows for comparison the same template as the red line but at $z=1.42$ as in the case the detected line was the CaT 8542\\AA\\ line. In the up left corner, a zoom view of the identified CaT(8662\\AA) absorption line is shown together with its best fit modeling ($\\sigma_{v}$=340 km/sec, red solid line) and other two templates corresponding to $\\sigma_{v}$=280 and $\\sigma_{v}$=460 km/sec. In the bottom box, the observed sky emission is reported for comparison.} \\end{figure} ",
        "conclusions": "We obtained a medium resolution (R=2000) LBT-LUCIFER spectrum of the compact early-type galaxy S2F1-142 at $z\\simeq1.4$. In the observed wavelength range, between 2.0 $\\mu$m and 2.3 $\\mu$m, we detected a single absorption line at 2.0668 $\\mu m$ that has been identified as the third CaT line at $8662$ \\AA. By means of this line we measured $z=1.386 \\pm 0.001$ and we estimated $\\sigma_{v}=340 ^{-60}-{+120}$ km/s. The corresponding virial mass has been calculated as 3.9$\\times10^{11}$ M$_\\odot$, compatible with the stellar mass estimates obtained assuming IMFs less dwarf reach than a Salpeter one.  \\noindent In Figure 7 (left panel) our target galaxy S2F1-142 is located on the size-mass diagram (filled magenta circle), together with the distribution of ETGs in the local Universe (gray circles,  \\citealt{me13}, \\citealt{ber13}) and other up to now available measures for galaxies at $z\\ga 1$ (blue symbols). All the mass estimates have been determined on the basis of a \\citet{c3} IMF. S2F1-142 is a compact galaxy with an effective radius more than 3 times smaller than expected from the average local value of galaxies with the same stellar mass content, i.e. $<R_{e}>(3.5\\times 10^{11}$ M$_\\odot$) = 10.2 kpc. In the right panel, the local distribution of the velocity dispersion values of ETGs in the local universe is reported together with the average values (solid line) and with the same sample of high redshift galaxies reported in the left panel. It is noticeable the large scatter of both the relations. First of all, we try to look for a local counterpart of S2F1-142, that is a local galaxy with the same mass, effective radius and velocity dispersion. In the sample of \\citet{ber13}, we find about 20 galaxies with stellar mass and $R_{e}$ compatible with those of S2F1-142 (i.e., $\\mathcal{M}_{stars}$ within 2 and 6 $\\times10^{11}$ M$_\\odot$ and $R_{e}$ between 2.5 and 3.5 kpc). Their velocity dispersion values are within 170 and 380 km/s, thus compatible within 1 $\\sigma$ error with $\\sigma_{v(re/8)}$. This means that S2F1-142 can be considered an extreme object but still similar to other extreme objects populating the local universe. Finally, we can compare S2F1-142 with other high redshift galaxies. With respect to the galaxy presented in \\citet{ono}, S2F1-142 has a a similar stellar mass content (i.e. $\\mathcal{M}_{* ONO} $ is 2.6x10$^{11}$M$_\\odot$ that is  0.6 times $\\mathcal{M}_{*}$ of S2F1-142), while its effective radius is 1.9 times smaller than $R_{e\\ ONO}$, where $R_{e\\ ONO}=5.79$ kpc. In case both the two galaxies were virialized, we than expect that S2F1-142 has a velocity dispersion $\\sigma_{v}=1.4 \\ \\sigma_{v\\ ONO}$, that is a value of its velocity dispersion of 378 km/sec, compatible with the measured $\\sigma_{v}$=340$^{-60}_{+120}$ km/sec.  We obtain a similar result even if we consider the 2 galaxies from \\citet{ca09} which have a larger difference in mass with respect to the previous case (i.e. $M_{*\\ CAP}=3.5-5.1 x 10^11 $M$_\\odot$, that are 0.08-0.12 times the stellar mass of S2F1-142) but similar effective radii (i.e. $R_{e\\ CAP}=2.16-1.81$ that are 0.7-0.6 times the effective radius of S2F1-142). In this case, the expected value of $\\sigma_{v}$ for S2F1-142 is 2.9-2.2 the values of the two galaxies of \\citet{ca09}, that is $\\simeq$ 330 km/sec, again compatible with the measured valued  $\\sigma_{v}$=340$^{-60}_{+120}$ km/sec.    \\vskip 0.5 truecm As a final test, we check if those galaxies which have a quite lower value of $R_{e}$ than expected from their stellar mass content and from the {\\it average} local size-mass relation (like S2F1-142) have a velocity dispersion correspondingly higher with respect to the {\\it average} sigma-mass relation, where the correspondence is evaluated in the hypothesis of virialization.  In fact, in a virialized system, it is expected that $\\sigma_{v}$, $R_{e}$ and M$_{star}$ are related in the following way:  $$M_{vir}=f M_{star}=c \\frac{R_{e} \\sigma^{2}}{G} $$  \\noindent where $c$ is the virial coefficient (e.g., \\citealt{ca06}) and $f$ is the ratio between  the total virial mass and the stellar mass. Assuming $\\frac{f}{c}$  constant, we expect that the velocity dispersion of two galaxies with different stellar mass content M$_{star 1}$ and M$_{star 2}$, and different effective radii $R_{e 1}$ and $R_{e 2}$ relate themselves as: $$\\frac{\\sigma_{1}}{\\sigma_{2}} =\\sqrt{ \\frac{M_{star 1}}{R_{e 1}}  \\frac{R_{e 2}}{M_{star 2}}}$$  \\noindent Thus, a galaxy like S2F1-142, whose $R_{e}$ is more than three times the average local value, is expected to have a velocity dispersion almost twice the mean local value of ETGs with its stellar mass. From the relation in Figure 7 (right panel), the local value of $\\sigma_{v}$ for galaxies with stellar mass equal to 3.5$\\times10^{11}$ M$_\\odot$ is 250 km/s (within a circular aperture of $R_{e}$/8), that is only 1.5 times smaller than $\\sigma_{v(re/8)}$ of S2F1-142-like galaxies  (1.7 considering 1$\\sigma$ error). To further check this evidence, in Figure 8, we report $\\sigma_{v}/\\sigma_{v\\ z=0}$  versus $R_{e}/R_{e\\ z=0}$, where the values of $\\sigma_{v\\ z=0}$ and $R_{e\\ z=0}$ are the average local values (green lines in Figure 7, left and right panel, respectively) for galaxies with the same stellar mass content. Symbols are the same as in previous Figure 7. It is immediately noticeable that the distribution of these quantities does not appear to depend on the redshift. The best fitting of the whole set of data results to be  $$\\frac{\\sigma_{1}}{\\sigma_{2}} =[ \\frac{M_{star 1}}{R_{e 1}}  \\frac{R_{e 2}}{M_{star 2}}]^{\\alpha}$$  \\noindent where $\\alpha$ is 0.20, and not 0.50 as expected on the basis of the previous hypothesis. The easiest way to explain this finding is to release the hypothesis that $\\frac{c}{f}$ is constant for all the ETGs, assuming on the contrary that it depends on their compactness. If this would be the case, the result $\\alpha=0.2$ translates into $\\frac{f}{c} \\sim [\\frac{M_{star}}{R_{e}}]^{-0.60}$.   \\noindent As pure speculation, we make the hypothesis that the virial coefficient $c$ is not systematically dependent on the compactness of galaxies. The above finding can then be read as if at fixed mass the dark matter fraction $f$ is higher in more diffuse galaxies while it is smaller in the most compact ones. Compact galaxies, both at high redshift and in the local universe, could have a lower dark matter content with respect to their more diffuse counterparts with the same stellar mass. This evidence would be consistent with the dry mergers scenario, where the more relaxed galaxies are the product of stars stripping from low mass satellites by more compact objects, which assemble the stellar particles of the satellite at large radii in halo-dominated regions of the massive host. This process strongly increases the size of the bulge into regions with higher dark matter fractions, leaving the inner host structure almost unchanged \\citet{hil12}. At the same time, this process does not seem an evolutionary effect, since the consequent relation between dark matter content and compactness is observed at any redshift.  \\vskip 0.3 truecm On the other hand, making the opposite hypothesis than before, that is the dark matter fraction in ETGs is constant and not dependent on their compactness, the previous finding would translate into ${c} \\sim [\\frac{M_{star}}{R_{e}}]^{+0.60}$. The virial coefficient $c$, at least for an isothermal sphere model, increases at increasing values on the profile Sersic fitting parameter $n$ (\\citealt{bert}), so that we can see $c$ almost proportional to $n$. This means that under the hypothesis of fixed dark matter fraction, the above finding translates into an expected dependence of the luminosity profile steepness on the compactness of the galaxies. This, differently from the previous case, would be in contrast with the inside-out dry mergers models, that foresees a larger value of $n$ for the less compact evolved ETGs (\\citealt{pat13}).  \\vskip 0.3 truecm   Finally, another simpler explanation of our finding of a high value of $\\alpha$ in  $\\frac{\\sigma_{1}}{\\sigma_{2}} =[ \\frac{M_{star 1}}{R_{e 1}} \\frac{R_{e 2}}{M_{star 2}}]^{\\alpha}$, without invoking varying dark matter content or non homology, could be that the stellar mass content of the apparently more compact galaxies is  overestimated, and their real compactness is thus less extreme than assumed. A possible reason of a systematic overestimate of stellar masses of the more compact galaxies could be related to the IMF assumed to transform their luminosity into a mass estimate. If this would be the case, we can expect that many of the up to now claimed compact ETGs are actually more relaxed galaxies for which the stellar mass estimate has been overestimated by assuming a too steep IMF (i.e., bottom heavy IMF). It is worth to note that evidences of a non universal IMF have been recently found out by many authors (e.g., \\citealt{ca12}, \\citealt{con13}, \\citealt{sl13}) even if a common view of its possible variation with density is far to be reached. In particular, while \\citet{con13} find evidences that favor a steeper IMF in more massive ETGs, \\citet{sl13} demonstrated the necessity of a bottom light IMF in a particularly massive elliptical galaxy at $z=2.14$. Under this last hypothesis, our finding on the connection  between the average size-mass relation and the average sigma-mass relation suggests the possibility that the stellar mass of those galaxies which appear much denser than the average relation  is lower than deduced by assuming a universal IMF, thus better supporting the \\citealt{sl13} finding. It is anyhow worth to note that stellar mass estimates are often affected by large uncertainties not only related to the assumed IMF, but also to the uncertain age and metallicity of their stellar populations, and to the inaccuracy of the stellar populations models themselves.    \\vskip 0.5truecm Summarizing, S2F1-142 is a compact galaxy that is fully consistent with the already available measures of $R_{e}$ and $\\sigma_{v}$ both in the local and in the distant universe. Furthermore, both in the local and in the distant universe, we note evidence that the size-stellar mass relation and the $\\sigma$-stellar mass one cannot be simply related to each other by assuming the virial theorem with an average universal dark matter fraction and/or a strict homology of all the ETGs. The latter two quantities should depend on the galaxies compactness. Alternatively, many of the up to now claimed compact ETGs are actually more relaxed galaxies for which the stellar mass has been overestimated, for example because of the assumption of a too steep IMF (i.e., bottom heavy IMF), or alternatively, because of some stellar parameters systematic.  \\begin{figure} \\includegraphics[width=8.2cm]{PaperSpec_fig6bis2.ps} \\caption{For the same samples of previous Figure 7 (reported with the same symbols), the distribution of $\\sigma/\\sigma_{z=0}$ values is reported versus $R_{e}/R_{e\\ z=0}$ values, where $\\sigma_{z=0}$ and $R_{e\\ z=0}$ are the average local values as derived from the relations in Figure 7 (green lines). } \\end{figure}"
    },
    "1402/1402.1977_arXiv.txt": {
        "abstract": "{We have revisited the series of observations recently used to infer a temporal variation of the interstellar helium flow over the last forty years. Concerning the recent  IBEX-Lo direct detection of Helium neutrals, there are two types of precise and unambiguous measurements which do not rely on the exact response of the instrument: the count rate maxima as a function of the spin angle, which determines the ecliptic latitude of the flow, and the count rate maxima as a function of IBEX longitude, which determines a tight relationship between the ecliptic longitude of the flow and its velocity far from the Sun. These measurements provide parameters (and couples of parameters in the second case) remarkably similar to the \\textit{canonical, old} values. In contrast, the preferential choice of a lower velocity and higher longitude reported before from IBEX data is based only on the count rate variation (at each spin phase maximum) as a function of the satellite longitude, when drifting across the region of high fluxes. We have examined the consequences of dead time counting effects, and conclude that their inclusion at a realistic level is sufficient to reconcile the data with the \\textit{old} parameters, calling for further investigations. We discuss the analyses of the STEREO pickup ion data and argue that the statistical method that has been preferred to infer the neutral flow longitude (instead of the more direct method based on the pickup ion maximum flux directions), is not appropriate. Moreover, transport effects may have been significant at the very weak solar activity level of 2007-2009, in which case the longitudes of the pickup ion maxima are only upper limits on the flow longitude. Finally, we found that the use of some flow longitude determinations based on UV glow data are not adequate. Based on this global study, and at variance with recent conclusions we find no evidence for a temporal variability of the interstellar helium flow. This has implications for inner and outer heliosphere studies.} ",
        "introduction": "The low energy, neutral helium flow measured in the solar system has its origin in the $\\simeq$25 km.s$^{-1}$  relative motion of the Sun with respect to the ambient interstellar cloud. Interstellar helium measurements are crucial because, at variance with hydrogen, helium atoms experience very little charge-exchange with the solar wind ions, and as a consequence the helium flow velocity vector is the best indicator of the actual relative motion. The precise direction of the helium flow also serves as a reference for measurements of other species, and inferred differences can be used to derive the detailed coupling between the interstellar and solar plasmas (e.g. \\cite{izmo99,lall05}. Possessing a very accurate value of the interstellar flow parameters is also particularly timely, because the two Voyager spacecraft are currently producing new and exciting data on the structure of the inner and outer heliosheath, including the recent crossing of the heliopause by Voyager 1 \\citep{gurnett13,stone13,krimigis13}.  After the pioneering measurements of the HeI 58.4nm solar  backscattered radiation \\citep{weller74,ajello78,ajello79,weller79,weller81}, the first  \\textit{in situ} measurements of interstellar helium were obtained, starting with pickup ions \\citep{mobius85,mobius95} and then with direct neutral helium atom detection \\citep{witte93}. Extensive efforts were done in 2003-2004 (ISSI workshop) to analyze in a synthesized manner all available datasets. In particular, the new analyses resolved the remaining strong discrepancies between flow temperature and velocities that were independently derived from UV and particle data \\citep{lall04a,lalluvcs}. A canonical set of parameters, i.e. flow direction, flow speed, and temperature, was presented by \\cite{mobius04} based on a synthesis of all data. Those flow parameters are very close to the parameters of \\cite{witte04} derived from the Ulysses atom detection data.  More recently, the Interstellar Boundary EXplorer (IBEX) mission has produced unprecedented measurements of neutral atoms in the heliosphere, in particular energetic atoms generated in the outer heliosphere \\citep{mccomasribbon} as well as low energy atoms of direct interstellar origin, especially helium atoms \\citep{hlond12}. The analysis of the first two years  of IBEX-Lo helium data (2009-2010) led to the derivation of helium flow parameters quite different  from the \\textit{canonical} values, in particular a flow longitude of 79.2$^{\\circ}$, instead of 75.4$^{\\circ}$, and a speed of 22.8 km.s$^{-1}$ instead of 26.3$\\pm$0.5 km.s$^{-1}$  \\citep{bzozo12,mobius12}. Soon after those new parameters were published, \\cite{drews12} announced similar values of the flow longitude, based on their analysis of pickup ions (PUIs) detected by STEREO/PLASTIC.  The value of the flow velocity has a strong influence on the heliospheric interface structure, and indeed the IBEX-Lo low speed was used as an argument to dismiss the existence of an interstellar  bow-shock ahead of the heliopause \\citep{mccomasbowshock}. It also motivated new theoretical studies about the actual nature of the outer heliosheath \\citep{zank13,zieger13}. A compilation of all results obtained since the 70's and including the new IBEX-Lo and STEREO determinations has been interpreted as a statistically significant demonstration of a temporal change in the longitude of the interstellar helium flow velocity vector over the last decades \\cite{crim13}, which is very surprising owing to the large mean free paths against collisions in low density clouds such as the circumsolar one.  In this work we revisit the most recent datasets on the helium flow and their interpretations. Section 2 is devoted to the IBEX-Lo results that have a very strong statistical weight in the analysis of the longitude temporal variation, due to the accuracy of the measurements. In section 3 we discuss the other data, i.e. PUIs and backscattered radiation measurements. Section 4 synthesizes the data and most reliable parameters and discusses the likelihood of temporal variations of the flow over the last decades.  \\begin{figure*} \\centering \\includegraphics[width=\\linewidth]{curvature.jpg} \\caption{Top: An extremely solid measurement of IBEX-Lo is the Earth longitude of the maximum neutral atom flux $\\lambda_{E}$, where the entrance field (hatched cones) is directed against the flow ($\\simeq$ ortho-radial bulk flow). Each value of $\\lambda_{E}$ corresponds to a continuum of pairs of flow longitude and speed values, as illustrated by the two (black and red) trajectories. Bottom left: the corresponding longitude-speed relationship inferred by \\cite{mobius12} (black line) and based on the measurement of $\\lambda_{E}$ with IBEX-Lo. Also shown is the central part of the chi-square parametric image of \\cite{bzozo12}, that also corresponds to the same longitude-speed link. The \\textit{consensus} pair of values is fully compatible with those relationships. Bottom right: the temperature-longitude derived by \\cite{mobius12} based on the flux measurements during the rotation around the IBEX spin axis. The authors considered both the absence (top curve) and the presence (bottom curve) of high-rate suppression of events (or dead-time). The \\textit{consensus} temperature of $\\simeq$ 6,300K is compatible with the \\textit{consensus} flow longitude of $\\simeq$ 75.4 $^{\\circ}$ for a dead time on the order of 5ms or slightly higher, i.e. realistic values according to \\cite{mobius12}.} \\label{figcurve}% \\end{figure*} ",
        "conclusions": "We have revisited the series of measurements used by \\cite{crim13} to infer a temporal increase of the helium flow longitude over the last decades. Motivated by the fact that both the flow latitude and the combination of speed and longitude deduced from IBEX-Lo count rate maxima are very surprisingly fully compatible with the canonical set of parameters derived in 2004, and that the preferential choice of a lower speed and higher longitude depends essentially on the count rate response to the atom flux, we have examined the potential consequences of high rate suppression of events (dead time effects) likely present in the IBEX data \\citep{mobius12}. The study shows that the \\textit{old} set of parameters, including a flow longitude of 75$^{\\circ}$, is as likely as the \\textit{new} one, provided a dead time of the order of 7ms is included, i.e. at a level estimated to be realistic by these authors. We conclude that the IBEX-Lo parameters can not be used for a temporal study until such effects are further examined. Based on the two mentioned coincidences it is very likely that the \\textit{old} set of parameters (and even probably a slightly lower longitude), is indeed the actual solution.  Based on published models \\citep{chalov06} we argue that PUI measurements can provide only an upper limit on the flow longitude, due to transport effects. At variance with ACE data in 2000 \\citep{gloeckler04}, the STEREO data were recorded at a time of particularly low activity, which favors the diffusion along the magnetic field lines, and potentially produces displacements of the PUI average trajectories by a up to 5$^{\\circ}$, according to \\cite{chalov06}. In this respect it is particularly interesting that the upwind maxima are systematically shifted to higher longitudes compared to the downwind focusing cones, as noted by \\cite{drews12}.  We also argue that the statistical study of the STEREO PLASTIC data is not adequate, being based on the assumption of dominant and smoothly varying ionization by UV photons. This assumption is highly questionable, since a large fraction of helium ions originates from electron impact ionization, especially within the Earth orbit, and this electron impact ionization is  a highly variable effect linked to dense solar wind streams and their suprathermal electron  tails \\citep{lalluvcs}.  We have also argued that the use of a fraction of the backscatter measurements and errors is inappropriate, in particular the flow direction from \\citep{ajello78,ajello79} that is derived from H Ly$\\alpha$ data and the \\textit{eye-fit} preliminary direction of \\cite{dalo84}. We also revisited the error bar for Nozomi Helium data \\citep{hiromu08}.  Fig \\ref{figtemporal} displays the various longitude measurements and corresponding errors we believe can be used at present, according to the above remarks. Longitudes derived from PUI data (augmented by 2 $\\sigma$) serve as upper limits only. Actually one PUI measurement, the 2000 ACE data \\citep{gloeckler04} is adding a constraint. A weighted linear fits to those data shows no signs of temporal increase (including or not the PUIs upper limits). This stability is in agreement with the absence of variation between the early value of the H flow direction (e.g. $\\lambda$=72$^{\\circ}$ from \\cite{ajello78}) and the most recent one with SOHO/SWAN \\citep{lall10}. Further analyses of the whole set of IBEX-Lo data at high temporal resolution, and taking into account as carefully as possible dead-time correction and electron noise should provide the most accurate parameters on the He flow among all experiments, allowing better inter-comparisons and determinations of temporal trends, if they do exist."
    },
    "1402/1402.4908_arXiv.txt": {
        "abstract": "{ We analyze the dependence of disk morphology (arm class, Hubble type, bar type) of nearby spiral galaxies on the galaxy environment by using local background density ($\\Sigma_{n}$), projected distance ($r_{p}$), and tidal index ($TI$) as measures of the environment. There is a strong dependence of arm class and Hubble type on the galaxy environment, while the bar type exhibits a weak dependence with a high frequency of SB galaxies in high density regions. Grand design fractions and early-type fractions increase with increasing $\\Sigma_{n}$, $1/r_{p}$, and $TI$, while fractions of flocculent spirals and late-type spirals decrease. Multiple-arm and intermediate-type spirals exhibit nearly constant fractions with weak trends similar to grand design and early-type spirals. While bar types show only a marginal dependence on $\\Sigma_{n}$, they show a fairly clear dependence on $r_{p}$ with a high frequency of SB galaxies at small $r_{p}$. The arm class also exhibits a stronger correlation with $r_{p}$ than $\\Sigma_{n}$ and $TI$, whereas the Hubble type exhibits similar correlations with $\\Sigma_{n}$ and $r_{p}$. This suggests that the arm class is mostly affected by the nearest neighbor while the Hubble type is affected by the local densities contributed by neighboring galaxies as well as the nearest neighbor. } ",
        "introduction": "The structure of a spiral galaxy is characterized by two primary components, disk and bulge of which radial luminosity distributions are well represented by an exponential profile \\citep{fre70} and $r^{1/4}$ law profile \\citep{dev53}, respectively. The ratio of the luminosities in these two components, the disk-to-bulge ratio, is a well defined quantity which can be used as a proxy of Hubble type. The spiral arms and bars are the secondary components displaying perturbations imposed on the disk. However, they are as important as the primary components, disk and bulge, because their structures are closely related to the internal and external properties of the galaxies \\citep{elm11} and they are thought to drive secular evolution of the spiral galaxies \\citep{kor79,pfe90,fri93,fri95,zha96,zha98,zha99, kor04}.  In Hubble's morphological classification of galaxies, spiral galaxies are divided into subtypes, known as Hubble stages (early-, intermediate- and late-type), according to the dominance of bulge luminosity, arm openness, and patchness of the spiral arms. There are quantitative measures of the first and second properties, namely, bulge-to-disk ratio and the pitch angle of the spiral arms. They are considered to reflect the physical properties of the disk morphology. But it is difficult to directly relate the patchness with the physical properties of galaxies because it is affected  by the distance of a galaxy. Thus, recent morphological studies of spiral galaxies consider the bulge-to-disk ratio and arm openness as the working criteria to distinguish the Hubble stage of spiral galaxies. Early-type spirals have tightly wound spiral arms and a large bulge-to-disk ratio, while the late-type spirals have loosely wound spiral arms and a small bulge-to-disk ratio.  However, the openness of the spiral arms is not the only property of spiral arm structures. \\citet{elm82,elm87} introduced arm classes, which are sometimes denoted as AC, to distinguish the arm structures on the basis of the degree of symmetry and continuity together with the number of arms. The 12 arm classes introduced by \\citet{elm82} were reduced to 10 arm classes, as two classes that are related to the presence of a bar (AC10) and companion (AC11) were removed. However, most previous studies divided the arm classes into two or three broad classes: (1) grand design (G) and flocculent (F) or (2) grand design (G), multiple-arm (M), and flocculent (F). For the case of two broad classes, multiple-arm is considered as a subclass of the grand design arms. Grand design arms have two long symmetric arms, while flocculent arms have a large number of chaotic fragmented arms. Multiple arms exhibit intermediate properties between the grand design and flocculent arms. Some multiple-arm galaxies have short grand design arms in the inner parts and multiple arms in the outer parts. Some of the outer arms in multiple-arm galaxies are likely to be connected to produce pseudo-rings \\citep{elm82}.  \\begin{figure*}[!t] \\center \\includegraphics[scale=0.90]{Fig1.eps} \\caption{Sample of spiral galaxies. Grand design galaxies, multiple-arm galaxies and flocculent galaxies are displayed in the top panels, middle panels and bottom panels, respectively. } \\end{figure*}   There have been many investigations regarding the relationships between arm classes and internal properties, including luminosity, color, disk-to-bulge ratio, and morpholgy. The variation in the mean luminosity and colors of galaxies across different arm classes is small~\\citep{ann13}. Conversely, considerable differences are observed in the galaxy size~\\citep{elm87,ann13} and  shape of rotation curves~\\citep{elm90,biv91}. There seems to be little or no dependence of the arm classes on the star formation rate~\\citep{elm86} and star formation related quantities such as the neutral hydrogen content \\citep{rom85}, CO surface brightness \\citep{sta87}, and supernova rate \\citep{mcc86}.  Arm classes exhibit some dependencies on the Hubble type \\citep{elm82,cho11, ann13}. Flocculent galaxies are more likely to occur in late-type galaxies, while multiple-arm galaxies are more frequent in early-type galaxies. Grand design galaxies are also more frequent in early-type galaxies, but the correlation is much weaker than multiple-arm galaxies. There is a correlation between the arm class and bar type (SA, SAB, and SB). The fraction of grand design galaxies increases from SA to SAB to SB, while flocculent galaxies show the opposite correlation. Multiple-arm galaxies exhibit nearly constant fractions regardless of the bar type. If we consider the Hubble type and bar type together, early-type SB galaxies are likely to have a high fraction of grand design galaxies; however, late-type SB galaxies exhibit no such tendency \\citep{elm89}.  The environment dependence of galaxy morphology is well understood for field galaxies \\citep{got03,par07} as well as cluster galaxies \\citep{dre80}. It is represented by the morphology-density relation (MDR) or morphology-radius relation (MRR). In other words, early-type galaxies (E/S0) are likely to be found in the dense environment or central regions of clusters, while late-type galaxies (Sp/Irr) are more frequently observed in the under-dense regions or outer parts of clusters. The environment dependencies of the disk morphology of spiral galaxies, which are represented by arm class, Hubble type, and bar type, is of interest because the arm class is closely related to the driving mechanism of the spiral structures \\citep{elm86,elm11}. Grand design arms are thought to be formed by density waves \\citep{lin64,ber89a,ber89b}, whereas flocculent arms are thought to be formed by stochastic self-propagated star formation \\citep{ger78, sei82}. Thus, understanding the environment of galaxies with different arm classes provides some information about the conditions for various modes of star formation.  \\begin{table} \\centering \\caption{Frequency of arm class, Hubble type and bar type of nearby spiral galaxies.} \\setlength{\\tabcolsep}{1.35mm} \\begin{tabular}{@{}cccccccccc@{}} \\hline & \\multicolumn{3}{c}{arm class}  & \\multicolumn{3}{c}{Hubble stage}  & \\multicolumn{3}{c}{bar type} \\\\ \\hline &  G &   M &    F &    e &   i &   l &   SA & SAB & SB \\\\ \\hline N   &  347 &  656 &  909 &  382 &  893 &  637 &  639 &  536 &  630 \\\\ \\hline \\end{tabular} \\end{table}   Early investigations into the environment dependence of arm classes \\citep{elm82,eed82} found that grand design galaxies are more frequent in high density regions, while flocculent galaxies favor low density regions. However, \\citet{giu89} reported that flocculent galaxies are more frequent in interacting or binary systems, which are likely to be present in dense environments. These contradictory claims may be due to poor statistics \\citep{giu94}. A recent study of the arm class and local background density of nearby spiral galaxies \\citep{ann13} confirmed the earlier findings of \\citet{elm82}, \\citet{eed82} and \\citet{elm87}. \\citet{ann13} found that there is a weak trend of increase in the fraction of grand design galaxies with increasing local background density, while a stronger opposite trend of a decreasing fraction with increasing density was found for flocculent galaxies. They derived the local background density from the projected distances to the 5th nearest neighbor. However, as described by \\citet{mul12}, the local background densities derived using the $n$th nearest neighbor method depend on two parameters $\\Delta V^{*}$ and $M^{*}$ which set the upper limit of the line of sight velocity difference between the target galaxy and its neighbor and the lower limit of the luminosity of the neighbor galaxies to be searched. Thus, it is necessary to choose a set of parameters ($n$, $\\Delta V^{*}$, $M^{*}$) suitable for the local background density and relevant to the formation of spiral structures.  The natural choice for $M^{*}$ appears to be the limiting magnitude that defines the sample galaxies as a volume-limited sample. However, other values of $M^{*}$ are also possible if the selected $M^{*}$ is relevant to a physical process being investigated. In case of $\\Delta V^{*}$, it would be different for galaxies in different environments because peculiar velocities depend on the large scale distribution of galaxies as well as nearby galaxies. Moreover, \\citet{par09} argued that $\\Delta V^{*}$ depends on the morphology of galaxies. They adopted $\\Delta V^{*}$=~400 and 600~km/s for late- and early-type galaxies, respectively, to search for the nearest neighbor galaxies on the basis of the pairwise velocity difference between the target galaxy and their neighbors \\citep{par08,par09}. Since we are dealing with galaxies in a variety of environments, a proper value  of $\\Delta V^{*}$ is yet to be determined. In addition to this, it is quite likely that disk morphology correlates well with other environment measures than local background density.  The purpose of the present study is to determine which measures of environment are closely related to the disk morphology. To do this, we first explore the dependence of local background densities on the neighbor search parameters, $M^{*}$, $\\Delta V^{*}$, and $n$, using the Sloan Digital Sky Survey (SDSS) Data Release 7 (DR7). Subsequently, we study local background densities of spiral galaxies, for which arm class, Hubble type, and bar type are known. We look for the appropriate set of $n$, $M^{*}$, and $\\Delta V^{*}$ that we could use to derive the local background densities relevant to the disk morphology.   In Section 2, we describe the observational data and method of arm classification. Local background density dependence of disk morphology with effects of neighbor search parameters on the local background densities are presented in Section 3. We descrive other measures of local environment in Section 4. Discussion on the relation between the disk morphology and the local environment is given in Section 5 and conclusions of the present study are given in Section 6 followed by acknowledgements in the last section. ",
        "conclusions": "It is apparent that the disk morphology is closely related with the local environment of galaxies. Arm class and Hubble type exhibit a clear dependence on three measures of the local environment: local background density ($\\Sigma$), projected distance to the nearest neighbor galaxy ($r_{p}$), and tidal index ($TI$). However, the correlation strength depends on the environmental measures used. Since measures of the galaxy environment are constrained by the neighbor search parameters, $M^{*}$ and $\\Delta V^{*}$, correlation strengths also depend on $M^{*}$ and $\\Delta V^{*}$. It was found that $\\Delta V^{*}$=~500~km/s is better than $\\Delta V^{*}$=~1000~km/s for all three measures, $\\Sigma$, $r_{p}$, and $TI$, but $M^{*}$ can be different for the measures of the selected environment.  Arm class exhibited the strongest correlation with $r_{p}$ when $r_{p}$ is derived from the neighbor search constraints $M^{*}=$~-18 and $\\Delta V^{*}$=~500~km/s. The fraction of grand design spirals decreases with increasing $r_{p}$, while that of flocculent spirals increases with it. Multiple-arm spirals exhibit a nearly constant fraction for $r_{p} <$ 1 $~h^{-1}$Mpc and decreases as $r_{p}$ increases further. $TI$ as well as $\\Sigma_{n}$ with $n=3$ and $5$ are also closely related with the arm classes of spiral galaxies; however, they exhibit a weaker correlation than $r_{p}$. It is worth noting that the luminosity constraint of $r_{p}$ that gives strongest correlation with arm class is $\\sim2$ magnitude brighter than the limiting magnitude that defines the present volume limited sample,  $M^{*}$=~-16.1. This means that the arm class is mostly affected by a massive neighbor galaxy, which is close enough to significantly perturb the disk instability of target galaxy.  The local background density, $\\Sigma_{n}$, is most closely related with the Hubble stage of spiral galaxies; however, $r_{p}$ and $TI$ also exhibit fairly good correlations with Hubble type. The best correlation was found for $\\Sigma_{5}$ with $M^{*}=-16.1$ and $\\Delta V^{*}$=~500~km/s. The fraction of early-type spirals increases with increasing $\\Sigma$ and $TI$ as well as decreasing $r_{p}$, while that of late-type spirals exhibited opposite trends. Intermediate-type spirals exhibited a nearly constant fraction along $\\Sigma$, $TI$, and $r_{p}$. The luminosity constraint of $r_{p}$ was $M^{*}$=~-18, while that for $\\Sigma$ and $TI$ was $M^{*}$=-16.  The bar type exhibited a weak dependence on the three environment measures, but there was a high frequency of strongly barred galaxies at high density regions where we expect relatively strong tidal strength and small separations of neighbor galaxies. The opposite is true for non-barred galaxies. Thus, the present results support the idea that bar formation is affected by the environment of galaxies.  The strong correlations between the frequencies of Hubble type of spiral galaxies, i.e., Hubble stage, and the local background density appear to originate from the MDR and MLR of galaxies, which govern the galaxy morphology throughout the galaxy formation process. This is the reason why the Hubble type exhibits a fairly strong correlation with $\\Sigma$. The strong dependence of the frequencies of the arm class on the projected distance to the nearest neighbor with the neighbor luminosity brighter than $M^{*}$=~-18 suggests that spiral arm structures, which are distinguished by arm class, are most affected by local perturbations caused by the nearest massive neighbor galaxy. If the perturbation is strong enough to drive the global density wave in the disk, grand design spiral arms are formed. However, if the perturbations are too strong, they break the symmetry of spiral arms and produce flocculent arms such as a single massive spiral arm."
    },
    "1402/1402.4139_arXiv.txt": {
        "abstract": "We have analysed the growth of Brightest Group Galaxies and Brightest Cluster Galaxies (BGGs/BCGs) over the last 3 billion years using a large sample of 883 galaxies from the Galaxy And Mass Assembly Survey. By comparing the stellar mass of BGGs and BCGs in groups and clusters of similar dynamical masses, we find no significant growth between redshift $z=0.27$ and $z=0.09$. We also examine the number of BGGs/BCGs that have line emission, finding that approximately 65 per cent of BGGs/BCGs show H$\\alpha$ in emission. From the galaxies where the necessary spectroscopic lines were accurately recovered (54 per cent of the sample), we find that half of this (i.e. 27 per cent of the sample) harbour on-going star formation with rates up to $10\\,$M$_{\\odot}$yr$^{-1}$, and the other half (i.e. 27 per cent of the sample) have an active nucleus (AGN) at the centre. BGGs are more likely to have ongoing star formation, while BCGs show a higher fraction of AGN activity.  By examining the position of the BGGs/BCGs with respect to their host dark matter halo we find that around 13 per cent of them do not lie at the centre of the dark matter halo. This could be an indicator of recent cluster-cluster mergers. We conclude that BGGs and BCGs acquired their stellar mass rapidly at higher redshifts as predicted by semi-analytic models, mildly slowing down at low redshifts. ",
        "introduction": "In the hierarchical model of structure formation, galaxies grow in size and stellar mass by accreting other galaxies and material from their surroundings. The brightest group and cluster galaxies (hereafter BGGs and BCGs) are the most extreme examples of this process and the most luminous objects known at the present epoch. Their properties are shown to be different to other early-type galaxies in clusters and groups \\citep[e.g. higher velocity dispersions than any other elliptical of the same mass, extended light profiles, and systematically brighter than what would be inferred from the luminosity function;][]{LOH06,LINDEN07,SHEN13}. The uniqueness of their properties has been attributed to their privileged location at the centre of their host group or cluster \\citep{HAUSMAN78}. However, their properties also correlate with the mass of their host cluster \\citep[e.g.][]{COLLINS98,BURKE00,BROUGH02}. \\\\\\\\ Most contemporary models of galaxy formation are based on the hierarchical assembly of dark matter halos \\citep[][]{TOOMRE77,WHITE78} in the $\\Lambda$CDM cosmology. In this paradigm, galaxies form at the centre of the halos. While N-body models accurately describe how the dark matter halos evolve, we are unable to simulate in detail the processes that lead to galaxy evolution. Halo abundance-matching (dark-matter-only) simulations and semi-analytical models (SAM) are two approaches that are commonly used in the literature. Halo abundance-matching models follow the behaviour of the $\\Lambda$CDM cosmology with gravity \\citep[e.g.][]{MOSTER13, LAPORTE13}. SAMs \\citep[e.g.][]{DELUCIA07, TONINI12} are a combination of N-body simulations and analytic descriptions of galaxy formation physics \\citep[i.e. star formation, dust extinction, AGN feedback, etc; for a review see][]{MUTCH13}. BCGs are particularly difficult to reproduce using these models, with their photometric colours tending to be bluer compared to observations, and their masses over-estimated \\citep{BOWER06, CROTON06, DELUCIA07,COLLINS09,SILK12}. \\\\\\\\ There are few models that focus solely on BGGs and BCGs. Among them we find \\citet{DELUCIA07} who used the Millennium N-body simulation of \\citet{SPRINGEL05} to model the development of BCGs over cosmic time. A similar  approach was used by \\citet{TONINI12}.  More recently, \\citet{MOSTER13} introduced an abundance-matching simulation of galaxy groups and clusters, using a statistical model constrained by observations. However, these simulations have not yet converged with observations. Despite these efforts, the assembly history and evolution of BGGs and BCGs is still poorly understood. SAMs and abundance-matching simulations predict a factor of three increase in the BCG stellar mass since $z=1$. On the other hand, earlier observational studies implied a very different result, arguing that BCG stellar masses 9 billion years ago were not very different to their stellar masses now \\citep{ARAGON98,BAUGH99,STOTT08,STOTT10,COLLINS09}. \\\\\\\\ Recently, \\citet{LIDMAN12} added new photometry from near-infrared imaging of clusters at $0.8<z<1.6$ to previous observations \\citep[][]{STOTT08,STOTT10}. Their results show that from $0.1<z<0.9$, BCGs grow in mass by a factor of $1.8\\pm0.3$. This is in closer agreement with the predictions from SAMs \\citep{DELUCIA07,TONINI12}. Also \\citet{LIN13}, with a sample from the \\textit{Spitzer} IRAC Shallow Cluster Survey  \\citep[ISCS;][]{EISENHARDT08} in the cluster mass range, $(2.4-4.5)\\times 10^{14}\\,$M$_{\\odot}$, found remarkably good agreement  with the SAM of \\citet{GUO11} over the redshift interval $0.5<z<1.5$ (growth by a factor of 2.3). However, below $z=0.5$, they found that the growth stalls, a result that is not seen in the models. In this paper we will examine this low redshift interval more closely ($0.09\\leq z \\leq 0.27$). \\\\\\\\ Another important prediction made by hierarchical formation models is that BCGs are assembling their mass through similar mass mergers with little gas present \\citep{DELUCIA07,LAPORTE13}. However, recent analyses have shown that some BCGs harbour on-going star formation \\citep{EDWARDS07,ODEA08,ODEA10,PIPINO09,LIU12,THOM12}. \\citet{LIU12}, with an optical sample selected from the Sloan Digital Sky Survey \\citep[SDSS][]{YORK00}, found that the star formation rate in BCGs is not always low, although it is not high enough to increase the stellar mass of the BCG by more than 1 per cent.  The existence of star formation in such old galaxies has important implications for simulations \\citep{TONINI12}. \\\\\\\\ It is also generally assumed that the central galaxy of a cluster is also the brightest and most massive one (BCG). However, \\citet[][]{BEERS83}, \\citet[][]{POSTMAN95}, \\citet[][]{LIN04}, \\citet[][]{PIMBBLET06}, \\citet{LINDEN07}, \\citet[][]{BILDFELL08}, \\citet[][]{HWANG08}, \\citet[][]{SANDERSON09}, \\citet[][]{COZIOL09}, and \\citet[][]{SKIBBA11}  have demonstrated that this is not always the case. The proposed explanation for this is ongoing halo merging. Observationally, this can be understood as different stages in the hierarchical clustering process \\citep{PIMBBLET08,BROUGH08}. \\\\\\\\ In this paper we analyse a sample of BGGs/BCGs selected from the Galaxy And Mass Assembly survey \\citep[GAMA; ][]{DRIVER11}. Our sample is one of the largest accessible to date, covering a wide range of halo mass ($10^{12}-10^{15}\\,$M$_{\\odot}$). Throughout this paper we separate groups from clusters by a halo mass cut of M$_{\\rm halo}=10^{14}\\,$M$_{\\odot}$\\footnote{We tested our analysis by separating groups from clusters through multiplicity, and we find the results to be consistent.}. We analyse the BGG and BCG stellar mass growth spanning $0.27\\leq z \\leq 0.09$, and compare our results with the galaxy formation and evolution models of \\citet{DELUCIA07}, \\citet{TONINI12}, and \\citet{MOSTER13}. In addition, we analyse the active galactic nucleus (AGN) and star formation (SF) activity within these galaxies as well as their position in their host halo. We are interested in: (a) exploring the impact of ongoing star formation on the growth of these giant galaxies, (b) comparing the properties of BGGs/BCGs that are not at the centre of the potential well with those that are, and (c) looking for correlations between the properties of BGGs/BCGs and the properties of the dark matter halos in which they live. Together, these will give us detailed information of the evolution history of group/clusters, and their brightest galaxies. \\\\\\\\ In Section 2 we describe the GAMA \\citep{DRIVER11} survey and the different catalogues used in our analysis. In Section 3 we revisit the M$_*-$M$_{\\rm halo}$ relationship. In Section 4 we describe our method for estimating the growth in BGGs/BCGs in the last 3 billion years. In Section 5 we analyse the BGG/BCG's AGN and Star formation activity. Section 6 examines the position of the BGGs/BCGs within their host halo. In Section 7 we discuss the main results obtained from our study, comparing them with model predictions. We present a summary of our final conclusions in Section 8. The cosmology adopted throughout this paper is H$_0$=70 kms$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_{\\Lambda}$=0.7. ",
        "conclusions": "We have analysed a large (883 galaxies) and homogeneous sample of low redshift ($0.09<z<0.27$) brightest group and cluster galaxies from the Galaxy And Mass Assembly (GAMA) survey. We summarise our conclusions below. \\\\\\\\(a) By comparing the BGG/BCG stellar mass in like-with-like halos we find no significant growth over this period of cosmic time. After comparing our results with previous analyses we conclude that BCGs acquire their stellar mass rapidly at early epochs ($z>0.3$). Below redshift $z\\sim0.3$ the stellar masses increase more slowly. This is possibly because the timescales or efficiencies for merging evolve. While observations are more consistent with the models (SAMs and abundance-matching models) at higher redshift, there are still small discrepancies at low redshifts. We stress the importance of taking into account the stellar mass~-~halo mass relationship for such a comparisons. \\\\\\\\ (b) We find that BGGs/BCGs are not completely dormant; at least 27 per cent of our sample host AGN and another 27 per cent are star forming. Their star formation rate decreases with stellar mass, from $10\\,$M$_{\\odot}\\,$yr$^{-1}$ at M$_*\\sim 10^{10.8}\\,$M$_{\\odot}$ to less than 1 M$_{\\odot}\\,$yr$^{-1}$ at M$_*\\sim 10^{11.6}\\,$M$_{\\odot}$. Therefore, BGGs are actively star-forming while BCGs are mostly quenched with higher fractions of AGN activity.  At stellar masses 10$^{11}\\,$M$_{\\odot}<\\,$M$_*<10^{11.4}\\,$M$_{\\odot}$  we find a bimodal population of star-forming and quenched systems. We conclude that despite the presence of star-formation in BCGs the SFRs are not high enough for star formation to contribute significantly to the stellar mass growth of these galaxies. \\\\\\\\ (c) We also examine the position of the BGGs/BCGs with respect of their dark matter halo and find that around $\\sim13$ per cent of the BGGs/BCGs are not centrally located. The halo properties, dominance and relative overdensity, in non-central BGGs/BCGs halos suggest that these halos have undergone recent mergers. This is further proven by the overall stellar mass difference between central and non-central BGGs/BCGs. We suggest that non-central BGGs/BCGs were the central galaxies in a smaller system that fell into the current system not long ago. The fraction of AGNs and star forming galaxies is roughly the same for central and non-central BGGs/BCGs."
    },
    "1402/1402.2881_arXiv.txt": {
        "abstract": "{} {There is an ongoing debate whether the solar activity cycle is overlaid with a long-term decline that may lead to another grand minimum in the near future. We use the size, intensity, and magnetic field strength of sunspot umbrae to compare the present cycle 24 with the previous one.} { We used data of the Helioseismic and Magnetic Imager onboard the Solar Dynamics Observatory and selected all sunspots between May 2010 and October 2012, using one image per day. We created two subsets of this data with a manual tracking algorithm, both without duplication. One is containing each sunspot (910 umbrae within 488 spots) and was used to analyze the distribution of umbral areas, selected with an automated thresholding method. The other one contains 205 fully evolved sunspots. We estimated the magnetic field and the total magnetic flux or those and discuss the relations between umbral size, minimum continuum intensity, maximum field strength and total magnetic flux. } { We find nonlinear relations between umbral minimum intensity and size and between maximum magnetic field strength and size. The field strength scales linear with the intensity and the umbral size scales roughly linear with the total magnetic flux, while the size and field strength level off with stronger flux. When separated in hemisphere and averaged temporally, the southern umbrae show a temporal increase in size and the northern umbrae stay constant. There is no temporal variation in the umbral mean intensity detectable. The probability density function of the umbral area in the ascending phase of the current solar cycle is similar to that of the last solar cycle. } { From our investigation of umbral area, magnetic field, magnetic flux and umbral intensity of the sunspots of the rising phase of cycle 24, we do not find a significant difference to the previous cycle, and hence no indication for a long-term decline of solar activity. } ",
        "introduction": "Sunspots are the most conspicuous manifestations of the solar magnetic activity in white light images, \\citep{bray_loughhead_64, wittamann_xu_1987}. They are at the top of the hierarchy of the photospheric magnetic structures, being about three orders of magnitude larger than magnetic flux tubes \\citep{keller_1992, wiehr_etal_2004, rouppe_etal_2005, solanki_etal_06,lagg_etal_2010}. Diameters of sunspot umbrae span a range between 1 to 20 Mm and their areas cover two orders of magnitude \\citep{solanki03r,rempel_schlichenmaier_2011,borrero_ichimoto_2011}. The umbrae of the largest sunspots are the coolest structures observed in the solar photosphere, their temperature lies some 2000\\,K below the quiet Sun photosphere \\citep{maltby_etal_86, kopp_rabin_92, balthasar_schmidt_1993, martinez_vazquez_1993, collados_etal_94}. The orientation of the magnetic field is close to the local vertical in umbrae \\citep{beckers_schroeter_1969,keppen_martinez_96,mathew_etal_2003, beck2008}.  The periodic variation of sunspot number, i.e., the solar cycle, was observed over 400 years \\citep{athay_warwick_1961, harvey_1992, spruit_2012, usoskin_2013}. It is traced with several quantities including the number and area of sunspots \\citep{hoyt_schatten_1996, hathaway_etal_2002, usoskin_2008, hathaway_2010}. The cyclic variation of umbral properties like intensity and magnetic field strength was subject of several investigations \\citep{livingston_02, penn_livi_06, penn_macdonald_07, penn_livi_10,  livingston_etal_2012, pevtsov_etal_2013}. Among others, \\cite{albregtsen78} and \\cite{albergtson_etal_1984} find that the umbral intensity fluctuates in phase with the solar cycle such that the umbrae are brighter at the end of a cycle. \\cite{norton_gilman_2004} find a decrease in the umbral intensity from early to mid phase of cycle 23 and an increase after the maximum of the cycle. In contrast, \\cite{mathew_etal_2007} report an invariant umbral intensity through the cycle 23. This is accompanied with a periodic variation of the maximum field strength of umbrae with the solar cycle phase \\citep{watson_etal_2011, pevtsov_etal_2011, reza_etal_2012a}.   The fraction of the umbral area to the sunspot area is apparently independent of the solar cycle. Using Greenwich photoheliographic results, \\cite{jensen_etal_1955}, \\cite{tandberg_1956} and \\cite{ringnes_1964} report a weak tendency for the average penumbra to umbra radius to change in phase with the solar cycle. \\cite{steinegger_etal_1990} and \\cite{brandt_etal_1990} find umbra-penumbra area ratios of 0.24 and 0.32 for small and large spots, respectively. \\cite{beck_chapman_1993} find the umbra-penumbra area ratio independent of the sunspot complexity and magnetic field strength and only weakly correlated with the umbral area. \\cite{mathew_etal_2007} find no variation of the umbral radius with the solar cycle. They note, however, a weak tendency of a secular trend in the northern hemisphere. \\cite{penn_macdonald_07} and \\cite{reza_etal_2012a} find no significant variation of the umbral area as a function of the solar cycle. \\citet[][hereafter SP]{schad_penn_10} observe a small variation in the umbral size correlated with the variation of the umbral intensity and the solar cycle.   The general view in old observations was that smaller sunspots have a higher continuum intensity and a lower magnetic field strength \\citep{bray_loughhead_64}. The underestimation of stray light in earlier works was noticed by \\cite{zwaan_1965}, who argued that much of the dependency of the continuum intensity on the sunspot size can be explained by stray light \\citep[see also][]{rossbach_schroeter_1970}. More recent investigations challenged the idea that the continuum-area dependency is an artifact of stray light \\citep{mcintosh_1981,stellmacher_wiehr_1988,martinez_vazquez_1993}. Most recent observations indicate that the continuum-area dependency is a real one and exists even after removal of stray light \\citep{mathew_etal_2007, wesolowski_etal_08, schad_penn_10, reza_etal_2012a}.    Measurements of umbral physical properties are prone to systematic and random errors. Intensity measurements of umbrae are affected by scattered light \\citep{mattig1971, martinez_pilet_1992, chae_etal_1998, beck_reza_2011a}. The field strength measurements in sunspot umbrae using the Zeeman splitting of spectral lines are contaminated on the one hand by molecular blends \\citep{wittmann_1972} and on the other hand by inaccurate atomic data \\citep{borrero_etal_2003}. The uncertainty in the measurements of the umbral area is due to the umbral fine structures \\citep{sobotka_etal93,schmidt_balthasar_94,lites_etal_04, rimmele_08, shimizu_etal_09}. Umbrae are more stable compared to penumbrae \\citep{robinson_boice_82,leka_skumanich_1998,rolf_etal_2010a, reza_etal_2012b, louis_etal_2012}. Larger spots tend to have darker umbrae, higher magnetic field strength, and longer lifetime compared to smaller spots \\citep[e.g.,][]{schrijver_1987}. However, there is a significant intrinsic scatter \\citep{lites_skumanich_1990,martinez_vazquez_1993,schlichenmaier_collados_2002}.   \\citet[][hereafter BOG]{bogdan_etal_1988} and SP study the umbral size distribution from 1917 to 1982, and during the solar cycle 23, respectively. They find that there is no significant temporal variation in the size distribution from one cycle to the other. Such an invariance in the umbral area distribution is noteworthy considering the drastic difference in the number of sunspots in different solar cycles. The distribution can be fitted with a lognormal function, motivated by the fragmentation and random processes relevant to sunspot formation \\citep[][and references therein]{bogdan_etal_1988}. The size distribution of sunspot groups in the Greenwich photoheliographic sunspot records was studied by \\cite{baumann_sol_2005} who find a lognormal distribution both for the instantaneous area (corresponding to BOG selection) and maximum area of sunspots.  The end product of a multiplicative and fragmentary process is a lognormal distribution \\citep{kolmogorov}. There are many applications of lognormal distributions in different branches of science \\citep[e.g.,][]{pauluhn_etal_2000, cme_lognormal, kobayashi_etal_2011}. In the photosphere, a lognormal distribution of the umbral area observed in about three orders of magnitude hints at a fragmentation origin for umbrae. BOG speculate that a lognormal distribution of emerging flux tubes indicates that they are probably the result of fragmentation of a large flux tube. As pointed out by \\cite{baumann_sol_2005}, it is unlikely to imagine that sunspots are purely the result of such fragmentation. It is perhaps more reasonable to imagine that they are  joint product of a fragmentation and coalescence of flux patches as discussed by \\citet{zwaan_1992} and \\citet{rolf_etal_2010a}.  Such a size distribution provides insights in the nature of the solar dynamo mechanism \\citep{parker_1979, ossendrij_2003}. Note however that most of the observed umbrae by BOG and SP were in their decay phase, so one should be careful about using their distribution as an argument for the origin of sunspot flux trunks in the convection zone. The fragmentation of the rising flux tube also happens near the solar surface \\citep{zwaan_1978, brandenburg_2005}. In a 16\\,Mm deep domain, \\cite{rempel_2011} finds fragmentation of the flux trunk down to the bottom boundary. This subsurface fragmentation causes fragmentation in the photosphere as observed by e.g., \\cite{louis_etal_2012}.    To distinguish between two sunspot formation scenarios (fragmentation or coalescence), one has to evaluate the distribution in an early stage of spot development, when the decay has not redistributed the surfaced flux. It would be interesting to evaluate the umbral area beneath the photosphere and prior to their photospheric appearance. If lognormal, then it hints at a fragmentation process in the convection zone. To get as close as possible to pre-emergence situation, we have evaluated the umbral area distribution in an early stage, when the decay has not redistributed the surfaced flux significantly (Sectoral.\\,\\ref{sec:selection}).  Magnetohydrodynamic simulations of emerging flux tubes \\citep{fan_2008}, active regions \\citep{cheung_etal10}, pores \\citep{cameron_etal_2007}, and sunspots \\citep{rempel_etal_09} became recently available. The flux emergence simulations model the rise of a twisted flux tube in the convection zone, its expansion, and fragments near the solar surface. In contrast in sunspot and pore simulations, the field lines are imposed on a relaxed hydrodynamic domain. Then, the dispersed flux elements partially evacuate and merge \\citep{kitiashvili_etal_2010} to form pores and sunspots. The coalescence of the emerged flux elements in the simulations is in agreement with observations of e.g., \\cite{bernasconi_etal_2002} and \\cite{rolf_etal_2010b}. The tethered-balloon model of \\cite{spruit_1981} predicts that the coalescence of small flux patches \\citep{pariat_etal_2004, reza_etal_2012a} is a surface phenomenon since these flux elements are deeply anchored. In other words, the emergence of large flux tubes does not fundamentally affect their integrity.  Due to the large range of parameters, simulations usually focus on some aspects of sunspot properties rather than addressing all issues. Umbrae in these simulations perform convective energy transport by umbral dots and transient light-bridges. \\citet[][his Fig.\\,11]{rempel_2011} presents a sunspot with a minimum (bolometric) umbral intensity of about 0.3\\,I$_{\\mathrm{c}}$. In another simulation he finds a minimum umbral intensity of about 0.22\\,I$_{\\mathrm{c}}$ \\citep[][his Fig.\\,5]{rempel_2012}. In both case, the unsigned flux of the simulated spot is about 10$^{22}$\\,Mx. Using the SIR code \\citep{sir92, sir_luis}, we performed spectral synthesis of the \\citet{rempel_2012} umbra and find a a minimum continuum intensity of 0.08\\,I$_{\\mathrm{c}}$ at a wavelength of 630\\,nm. This minimum intensity is close to what one expects for such a big sunspot (compare  with the sunspot of Nov 19, 2013).   \\cite{rempel_2011} presents a simulated sunspot in a 16\\,Mm deep domain. In this simulation, a considerable amount of subsurface magnetic field is fragmented due to convection. The photospheric manifestation of this fragmentation is the appearance of light-bridges or flux separation. He finds a deep-reaching outflow beneath the photosphere and down to the bottom boundary of the simulation domain. Such an axisymmetric outflow presumably supports the monolithic subsurface structure of sunspots \\citep{parker_1979a}. It is not clear what role the supergranulation flow plays in this process.  Sunspots decay  as a result of fragmentation \\citep{Petrovay_1997, valentin_2002}. The decay rate, which in turn determines the sunspot lifetime, is a function of anchoring depth and the corresponding convection timescales in such depths \\citep{moradi_etal_2010}. The life time of sunspots, $T$, falls in-between two extremes: $$\\frac{H_{\\mathrm{p}}}{\\it{v}_{\\mathrm{rms}}} \\ll T \\ll \\left(\\frac{r_0}{10\\,\\mathrm{Mm}}\\right)^2 \\times10\\,[\\mathrm{days}],$$ where $H_{\\mathrm{p}}$ is the pressure scale height, $\\it{v}_{\\mathrm{rms}}$ is a typical convection velocity, and $r_\\mathrm{0}$ the maximal radius of the spot \\citep{petrovay_moreno_1997, schlichenmaier_1999}.  \\cite{mcintosh_1981} reports on a long-lived sunspot in Aug 1966, which survived five solar rotations (137 days). While the convection timescale (the lower limit) is about 6\\,hr at a depth of 15\\,Mm, \\cite{rempel_2011} speculates that the lifetime of a sunspot is roughly an order of magnitude longer than the convection timescale (at the corresponding anchoring depth). The trend found in his simulation cannot be extrapolated to broader ranges since the aforementioned long-lived sunspot will need an anchoring depth of about four times the depth of the convection zone.  The prolonged duration of the minimum of the last solar cycle  (23) raises speculations about a lower activity level in solar cycle 24 and the possible approach of a new grand minimum \\citep{tripathy+etal2010,jain+etal2011,tap_etal_2011}. There are also indications that the population of small sunspots compared to the large ones in solar cycle 23 might be different compared to previous cycles \\citep{clette_lefevre_12}. This motivated us to study properties of umbrae in the current solar cycle (24) and compare it with previous cycles.    In this contribution, we measure the umbral area, evaluate the maximum field strength and minimum intensity of sunspot umbrae in solar cycle 24 based on the data obtained with the Helioseismic and Magnetic Imager \\citep[HMI,][]{scherrer_etal_2012, hmi_2012}, a filter instrument onboard the Solar Dynamics Observatory \\citep[SDO,][]{sdo}. From May 2010 till October 2012, 4229 sunspots were observed with some 6892 (910 unique)  umbrae. The criteria for the data selection are discussed in Sect.\\,2. We analyze the data and interpret the results in Sects.\\,3 and 4, respectively. A summary and conclusions are presented in Sect.\\,5. ",
        "conclusions": "We use HMI data to investigate properties of sunspots in the ascending phase of the solar cycle 24. We use one solar image per day and apply an automated thresholding method to measure the umbral area. The data was corrected for limb-darkening and foreshortening effects. We performed an inversion to retrieve the magnetic field vector for a hand-selected subset of 205 sunspots. We use the large sample (910 unique umbrae) to analyze the distribution of umbral sizes while the small sample enables us to study empirical relations and temporal variations of size, magnetic field strength, and continuum intensity. The relations between magnetic field strength, continuum intensity and size of the umbrae are similar to previous studies, an invariant size-field strength relationship is likely to exist.  The relation between umbral intensity and size shows a power-law behavior with an exponent similar that of the previous cycle. Compared to a linear fit, a power-law function fits better the dependency of intensity on magnetic field strength and size. The continuum intensity is a good proxy of the field strength while sunspot flux closely follows the umbral area. The umbral radius shows an increase in the southern hemisphere, while no significant temporal variation in the northern hemisphere was observed. The umbral intensity does not show a temporal variation in the ascending phase of cycle 24 similar to cycle 23. The umbral size distribution is also similar to the one in the rising phase of cycle 23. Using sunspot properties in ascending phase of cycle 24, we do not find any evidence for a significant decrease in solar activity as suggested by e.g., \\citet{livingston_etal_2012}. We will continue to monitor sunspot properties and their changes through the rest of this cycle, to examine the likelihood of a long-term decrease of the solar activity."
    },
    "1402/1402.2553_arXiv.txt": {
        "abstract": "We have used recent Very Long Baseline Interferometry (VLBI) observations by \\citeauthor{2013A&A...551A..97M} with a resolution of $\\approx 10$ mas to investigate the properties of faint sources selected from the Tenth Cambridge (10C) survey in the Lockman Hole. The 10C survey is complete to 0.5~mJy at 15.7~GHz and has a resolution of 30~arcsec. We have previously shown that this population is dominated by flat-spectrum sources below $\\approx 1$~mJy, in disagreement with several models of the faint, high-frequency sky. We find that 33 out of the 51 10C sources in the VLBI field (65 percent) are detected by the VLBI observations. The sources detected by the VLBI observations must have a high brightness temperature, thus ruling out the possibility that this faint, high frequency population is dominated by starbursting or starforming sources and indicating that they must be Active Galactic Nuclei. ",
        "introduction": "\\label{section:intro}  In \\citeauthor{2013MNRAS.429.2080W} (2013; hereafter Paper I) we studied a sample of 296 faint ($> 0.5$~mJy) sources selected from the Tenth Cambridge survey (10C; \\citealt{2011MNRAS.415.2708D,2011MNRAS.415.2699F}) at 15.7~GHz in the Lockman Hole. The 10C survey was made with the Arcminute Microkelvin Imager (AMI; \\citealt{2008MNRAS.391.1545Z}) and covers ten fields complete to 0.5~mJy, two of which are in the Lockman Hole. By matching the 10C catalogue to several lower-frequency surveys we investigated the radio spectral properties of the sources in this sample. We found a significant change in spectral index $\\alpha$  with flux density $S$ (where $S \\propto \\nu^{-\\alpha}$ for frequency $\\nu$) -- the median spectral index between 15.7~GHz and 610~MHz changes from $\\alpha = 0.75$ for 15.7-GHz flux densities $1.5 < S/ \\rm mJy < 10$ to $\\alpha = 0.08$ for lower flux densities ($0.3 < S/ \\rm mJy < 0.8$) (full details of this study are given in Paper I). This suggests that a population of faint, flat-spectrum sources is emerging at flux densities $\\lesssim 1 \\textrm{ mJy}$.  The properties of the 10C sample were compared with those of a sample of sources selected from the SKADS Simulated Sky (S$^3$; \\citealt{2008MNRAS.388.1335W,2010MNRAS.405..447W}) (full details are in Paper I).  We found that this simulation fails to reproduce the observed spectral index distribution and underpredicts the number of sources in the faintest flux density bin ($0.5 < S < 1~\\rm mJy$). We discussed two possible origins for this discrepancy. The first is that the source population at this flux density level is not dominated by Active Galactic Nuclei (AGN) in the form of Fanaroff and Riley type I (FRI; \\citealt{1974MNRAS.167P..31F}) sources, as predicted by the S$^3$ model, and is instead dominated by a population with flatter spectra, such as, for example, a population of starburst galaxies with unusually flat high-frequency spectra \\citep{2013ApJ...768....2M}. The second possibility is that FRI sources are not modelled correctly in the simulation and in fact themselves have much flatter spectra at high frequencies. In Paper I we suggested that the first option is unlikely, as it would require the number of starburst sources to be incorrect in the simulation by at least a factor of ten.  Information about the structure of the sources in a sample can provide useful information about their nature, meaning that high resolution studies are valuable when investigating the properties of a population. The angular size of sources in the high frequency (20~GHz) source population at much higher flux densities ($S > 40~\\rm mJy$) has recently been investigated by \\citet{2013MNRAS.tmp.1815C}. They used data from a 6~km baseline to split the sample into different populations and investigate their properties. They found that 77 percent of their sample are compact AGNs and 23 percent are extended AGN-powered sources. Here we study a sample of high frequency sources which are considerably fainter on even smaller angular scales.  In this paper we use recently published VLBI data (\\citealt{2013A&A...551A..97M}, hereafter M2013) to investigate further the nature of this faint 15.7-GHz source population. M2013 have recently conducted wide-field observations with the Very Long Baseline Array (VLBA) at 1.4~GHz of part of the Lockman Hole. The VLBA provides milliarcsec-scale resolution, so if a source is detected the emission must come from a very compact region and the brightness temperature must be very high ($\\gtrsim 10^6$ K). Any stellar non-thermal sources, such as supernova remnants, at redshift $\\gtrsim 0.1$ cannot have a surface brightness this high and thus any objects with $z \\gtrsim 0.1$ which are detected must originate in AGN.  In Section \\ref{section:methods} the different samples used in this paper are defined and the methods used to investigate the properties of these samples are described. The results and discussion are split into two parts -- the first part, Section \\ref{section:results1}, discusses the properties of the 10C sources in the VLBA survey area and the second part, Section \\ref{section:results2}, contains further insights into the nature of the M2013 VLBA sources. The conclusions are given in Section \\ref{section:conclusions}.  Throughout this paper the term `flat spectrum' refers to an object with spectral index $\\alpha \\leqslant 0.5$, `steep spectrum' to an object with $\\alpha > 0.5$ and `rising' to an object with $\\alpha < 0$. ",
        "conclusions": "\\label{section:conclusions}  Sixty-five percent (33/51) of the 10C sources in the VLBA survey area are detected by the VLBA observations, showing that these sources are AGN. The detected sources have a range of 15.7-GHz flux densities, with detected sources as faint as $S_{15.7~\\rm{GHz}} = 0.3$~mJy. These results rule out the possibility discussed in Paper I that the 10C population is dominated by starforming or starbursting sources and provides strong evidence for our conclusion that the faint, high-frequency population is dominated by AGN, such as FRI sources. The proportions of sources in the $\\rm S^3$ model are therefore probably roughly correct, i.e. the population at 1 mJy at 15~GHz is dominated by FRI sources. The spectral properties of these sources are not modelled correctly in the simulation, both because the flat spectrum cores are not correctly modelled, and because no spectral ageing has been included in the source model.  These results also show that there is a small but significant population (four out of 33) of very compact, steep spectrum sources in the 10C sample. They also show that in a number of cases the extended emission displayed by the sources in the VLBA sample is not caused by star formation as assumed by Middelberg et al., and is instead synchrotron emission produced by the AGN jets.  Multi-wavelength observations of these sources will be presented in a future paper and will provide further insights into the nature of this faint, high-frequency source population."
    },
    "1402/1402.2891.txt": {
        "abstract": "We demonstrate possible optical signatures of the Kehagias-Sfetsos naked singularity spacetimes representing spherically symmetric vacuum solution of the modified Ho\\v{r}ava gravity. In such spacetimes, accretion structures significantly different from those present in the standard black hole spacetimes occur due to the \"antigravity\" effect causing existence of an internal static sphere surrounded by Keplerian discs. We focus our attention on the optical effects related to the Keplerian accretion discs, constructing the optical appearance of the Keplerian discs, the spectral continuum due to their thermal radiation, and spectral profiled lines generated in the innermost parts of such discs. The KS naked singularity signature is strongly encoded in the characteristics of predicted optical effects, especially in the case of the spectral continuum and spectral lines profiled by the strong gravity of the spacetimes, due to the region of the vanishing of the angular velocity gradient influencing the effectivity of the viscosity mechanism. We can conclude that optical signatures of the Kehagias-Sfetsos naked singularities can be well distinguished from the signatures of the standard black holes. ",
        "introduction": "The Ho\\v{r}ava (or Ho\\v{r}ava-Lifshitz) gravity \\cite{Hor:2009:PHYSR4:} is one of the promising approaches to the quantum gravity that recently attracts strong attention. The Ho\\v{r}ava quantum gravity is a field-theory that is based on the ideas of the solid state physics \\cite{Lif:1941:ZhEkspTheorPhys:} and uses an anisotropic scaling of space and time. Its Lagrangian demonstrates the Lorentz invariance at low energies, while the Lorentz invariance is broken at high energies. In the Ho\\v{r}ava gravity the space dimensions are scaled as $x \\rightarrow bx$, while the time is scaled as $t \\rightarrow b^{z}t$ with $z$ being an integer; the Lorentz invariance is recovered for $z=1$ \\cite{Hor:2009:PHYSR4:,Hor:2009:PHYSRL:,Hor-MelT:2010:PHYSR4:,And-etal:2012:PHYSR4:,Gri-Hor-MelT:2012:JHEP:,Gri-Hor-MelT:2013:PHYSRL:,Ver-Sot:2012:PHYSR4:,Lib-Mac-Sot:2012:PHYSRL:}.  The solutions of the Ho\\v{r}ava-Lifshitz effective gravitational equations have been found in \\cite{Bar-Sot:2012:PHYSRL:,Bar-Sot:2013:PHYSRL:,Bar-Sot:2013:PHYSR4:}. The spherically symmetric solution, compatible with the Minkowski vacuum and having asymptotically the Schwarzschildian character, has been found in the framework of the modified Ho\\v{r}ava model and is given by the so called Kehagias-Sfetsos (KS) metric \\cite{Keh-Sfe:2009:PhysLetB:,Park Mu-In:2009:JHEP:}. The KS metric involves two parameters - the gravitational mass $M$ determining the distance scales in the metric, and the Ho\\v{r}ava parameter $\\omega$ reflecting the influence of the quantum effects. The character of the KS spacetime is governed by the dimensionless product $\\omega M^2$. Considering the Ho\\v{r}ava parameter to be an universal constant, the character of the spacetime is governed by the mass parameter $M$. If the product $\\omega M^2 \\geq 1/2$, the KS metric describes a black hole, while for mass parameter small enough, when $\\omega M^2 < 1/2$, there are no event horizons, and the metric describes a naked singularity. The observational limits on $\\omega$ presented in \\cite{Ior-Rug:2010:IJMPA:,Liu-etal:2011:GenRelGrav:,Ior-Rug:2011:IJMPD:} do not exclude the existence of stellar-mass KS naked singularities. Therefore, it is of crucial interest to look for signatures of KS naked singularity spacetimes that enable to clearly distinguish them from the black hole backgrounds due to the appearance of accretion phenomena.  For the accretion phenomena, the KS black hole spacetimes has been extensively studied in a series of works related both to the particle motion \\cite{Hak-etal:2010:MPLA:,Ali-Sen:2010:PHYSR4:,Abd-Ahm-Hak:2011:PHYSR4:,Ior-Rug:2011:IJMPD:,Eno-etal:2011:PHYSR4:,Hor-Ger-Ker:2011:PHYSR4:,Hak-Abd-Ahm:2013:PHYSR4:} and optical phenomena \\cite{Ama-Eir:2012:PHYSR4:,Eir-Sen:2012:PHYSR4:,Ata-Abd-Ahm:2013:AstrSpaScie:} that can be relevant for tests of validity of the Ho\\v{r}ava gravity. Quite recently, some works have been related to the KS naked singularity spacetimes. The nature of the geometrical properties of the KS spacetimes has been represented by the embedding diagrams in \\cite{Gol-etal:2013:submitted:}, the ultra-high-energy particle-collisions in the deep gravitational potential of KS naked singularity spacetimes have been addressed in \\cite{Stu-Sche-Abd:2013:submitted:}, and the circular geodesics were studied in \\cite{Vie-etal:2013:submitted:}, where close similarity to the properties of the circular geodesics of the Reissner-Nordstrom naked singularity spacetimes \\cite{Stu-Hle:2002:ActaPhysSlov:,Hac-etal:2008:PHRVD:,Pug-etal:2011:PHYSR4:} has been explicitly demonstrated. In all of the considered phenomena, an \"antigravity\" effect \\cite{Vie-etal:2013:submitted:} plays a fundamental role. The situations is similar also in the braneworld naked singularity spacetimes \\cite{Stu-Kot:2009:GenRelGrav:,Sche-Stu:2009:IJMPD:,Stu-Kol:2012:JCAP:,Ali-etal:2013:CLAQG:}, but it seems to be different from the Kerr naked singularity spacetimes \\cite{Stu:1980:BAC:, Vir-Ell:2002:PhRvD:,Stu-Sche:2010:CLAQG:,Stu-Hle-Tru:2011:CLAQG:,Pat-Jos:2011:CLAQG:,Stu-Sche:2012a:CLAQG:,Stu-Sche:2013:CLAQG:,Sche-Stu:2013:JCAP:,Kol-Stu:2013:PHYSR4:,Bam-Mal:2013:PhRvD:}.  Due to the \"antigravity\" effect, all the KS naked singularity spacetimes contain a \"static\" (or \"antigravity\") sphere where test particles can remain in stable equilibrium positions. The static sphere represents the innermost limit on existence of circular geodesics (corresponding to orbits with zero angular momentum). However, the properties of the circular motion strongly depend on the dimensionless parameter $\\omega M^2$ \\cite{Vie-etal:2013:submitted:,Stu-Sche-Abd:2013:submitted:}. For $0 < \\omega M^2 < (\\omega M^2)_{ms} = 0.2811$, stable circular orbits exist starting at the static sphere of radius $r_{stat}$ up to infinity. For $(\\omega M^2)_{ms} < \\omega M^2 < (\\omega M^2)_{ph} = 0.3849$, two distinct regions of stable circular geodesics exist - the inner one between the static radius $r_{stat}$ and the outer marginally stable orbit at $r_{OSCO}$, and the outer one extends between the inner marginally stable orbit $r_{ISCO}$ and infinity; at the region $r_{OSCO} < r < r_{ISCO}$ unstable circular geodesics are located. For $(\\omega M^2)_{ph} < \\omega M^2 < 1/2$, the inner region of stable circular orbits terminates at a stable photon circular orbit at radius $r_{ph/o}$, while an unstable photon circular orbit at $r_{ph/i}$ represents a limit of unstable circular geodesics located at $r_{ph/i} < r < r_{ISCO}$, located under the outer region of stable circular geodesics; at the region $r_{ph/o} < r < r_{ph/i}$, no circular geodesics are allowed \\cite{Vie-etal:2013:submitted:}.  The standard Keplerian accretion discs are allowed in the regions of the outer stable circular geodesics; in the case of $0 < \\omega M^2 < (\\omega M^2)_{ms} = 0.2811$, this is the whole region of the stable circular geodesics, in the other cases of the KS naked singularity spacetimes, such discs are limited by $r_{ISCO}$ from below. The inner Keplerian discs related to the inner region of the stable circular geodesics cannot be created by the standard accretion processes, but we can consider some other processes, as gravitational radiation of orbiting matter supplied by the standard (outer) Keplerian discs, with the angular momentum limited by $L=L_{ISCO}$ corresponding to the inner marginally stable circular orbit, or with larger values of angular momentum, $L > L_{ISCO}$, corresponding to matter inflowing from thick, toroidal accretion discs. The orbiting matter can succesively settle down to a nearly circular geodesic motion, forming and inner Keplerian disc.  An alternative possibility of creating an inner Keplerian disc occurs when particles falling freely from infinity enter high-energy collisions slightly above the static radius \\cite{Stu-Sche-Abd:2013:submitted:}. Some of the new particles resulting from the high-energy collisions can follow circular or nearly circular geodesics, being trapped in the \"antigravity\" region; particles on such circular geodesics can even have very large (covariant) specific energy (and specific angular momentum) since the high-energy particle collisions near the static radius can convert the rest energy of heavy infalling particles into kinetic energy of created light particles \\cite{Stu-Sche-Abd:2013:submitted:}.  Here we focus our attention on the optical phenomena related to the outer (standard) Keplerian discs orbiting the KS naked singularity spacetimes, and compare them to those related to the KS black hole spacetimes or Schwarzschild spacetimes. We discuss their appearance, their spectral continuum created by the thermaly radiating Keplerian discs, and, finally, we construct the spectral profiled lines generated by the innermost parts of the outer Keplerian accretion discs. For comparison, we take into account in some cases also the effects generated by the inner Keplerian discs or some parts of these discs. In our study, we assume the motion of photons and matter governed by the General Relativity laws, i.e., the Infra-Red end of the Ho\\v{r}ava theory, and possible effects of the Lorentz invariance violation, expected at the Ultra-Violet end of the theory \\cite{Hor:2009:PHYSR4:,Hor:2009:PHYSRL:,Keh-Sfe:2009:PhysLetB:,Park Mu-In:2009:JHEP:}, are not considered.  The KS geometry and the equations of motion of test particles and photons are given in Section 2. In Section 3, circular geodesics are discussed and properties of the photon motion in the KS naked singularity spacetimes are determined and compared to those corresponding to the KS black hole spacetimes (and Schwarzschild spacetimes). In Section 4, properties of the geodetical circular orbits in the KS naked singularity spacetimes are discussed and related to the Keplerian discs and the general accretion phenomena in these spacetimes. Then we present in Section 5 the appearance of the innermost parts of the Keplerian accretion discs under assumption of discs radiating at a given frequency, in order to demonstrate the effects of the gravitational lensing and the Doppler and gravitational frequency shift. In Section 6, we give the spectral continuum under assumption of thermaly radiating Keplerian discs (of the Page-Thorne type), and in Section 7 we construct the profiled spectral lines generated by the Keplerian discs if their radiation is assumed at a fixed frequency related, e.g., to some of the Fe spectral lines. The results of the calculations of the spectral effects are summarized and discussed in Section 8. Concluding remarks are presented in Section 9. ",
        "conclusions": "We have studied the optical phenomena related to the Keplerian accretion discs orbiting the KS naked singularity spacetimes. We have considered three optical phenomena: a) the appearance of the Keplerian discs demonstrating the influence of the strong gravity effects on the shape of the discs and the frequency shift of the radiation generated by them, b) the radiation flux of thermaly radiating (Page-Thorne) Keplerian discs, their temperature profile, and the corresponding spectral continuum observed at large distances, c) the spectral (fluorescent) lines generated by an outer irradiation of the Keplerian discs and profiled by the strong gravity effects of the spacetime.  The KS naked singularity spacetimes are separated into three distinct classes according to the properties of the geodetical (Keplerian) circular orbits in dependence on the dimensionless Ho\\v{r}ava parameter $\\omega < 1/2$. For all three classes of the KS naked singularity spacetimes, the innermost circular orbits correspond to the stable equilibrium point with $L=0$ at the so called static radius where particles are at rest relative to distant observers. The other properties of the circular geodesics are quite different in those three classes of the KS naked singularities.  For small values of the naked-singularity Ho\\v{r}ava parameter, $\\omega < \\omega_{ms}$, only stable circular orbits can exist at such KS naked singularity spacetimes -- no unstable circular orbits and no circular photon orbits exist. For mediate values of the Ho\\v{r}ava parameter, $\\omega_{ms} < \\omega < \\omega_{ph}$, two regions of stable circular orbits exist, being separated by a region of unstable circular geodesics -- no photon circular orbits are possible in such spacetimes. For large values of the naked-singularity Ho\\v{r}ava parameter, $\\omega_{ph} < \\omega < 1/2$, two regions of stable circular orbits exist again; the inner region is limited from above by a stable photon circular orbit; under the outer region of stable circular geodesic orbits a region of unstable circular geodesics extends being limited from below by an unstable photon circular orbit. No circular geodesics are allowed between the stable and unstable photon circular orbits.  In all three classes of the KS naked singularity spacetimes, there are regions where the gradient of the angular velocity of the circular geodesics is oppositely oriented, while in the black hole spacetimes, there is always $\\frac{d\\Omega_{K}}{dr} < 0$ along the whole Keplerian discs allowing thus for the MRI viscosity mechanism to work in whole the region of stable circular geodesics. In the KS naked singularity spacetimes with  $\\omega_{ph} < \\omega < 1/2$, the regions with $\\frac{d\\Omega_{K}}{dr} > 0$ and $\\frac{d\\Omega_{K}}{dr} < 0$ are separated, while in the spacetimes with $\\omega < \\omega_{ph}$, there is a radius where $\\frac{d\\Omega_{K}}{dr} = 0$ and these regions are smoothly matched there, but the MRI viscosity mechanism governing the standard Keplerian accretion stops its functioning there. We call the discs where $\\frac{d\\Omega_{K}}{dr} < 0$, having the standard MRI viscosity mechanism at work, the outer Keplerian discs, while those where a region with $\\frac{d\\Omega_{K}}{dr} > 0$ occurs we call the inner Keplerian discs. In the inner Keplerian discs we assume occurence of a non-standard viscosity mechanism creating a thermalized disc, or we assume that the accretion is governed by the gravitational radiation of the orbiting matter and the disc has to be cold. It should be stressed that a non-Keplerian accretion structures (toroidal or even spherical) are more probable to occur in the regions related to the inner Keplerian discs, nevertheles, occurence of limited inner Keplerian discs cannot be excluded and that is the reason why we have studied such structures in our paper.  All three considered optical phenomena, the appearance of the Keplerian discs, their spectral continuum and their profiled spectral lines demonstrate clear signatures of the KS naked singularity spacetimes, enabling us to distinguish them from the Keplerian discs orbiting the KS or Schwarzschild (Kerr) black holes. The strong, qualitative signatures of the KS naked singularity spacetimes are related mainly to their inner Keplerian discs (or inner accretion structures, if Keplerian discs are destroyed), and enable even strong restriction of the dimensionless parameter $\\omega$ in order to distinguish the three clasess of the KS naked singularity spacetimes. Moreover, even the optical phenomena related only to the outer Keplerian discs where the standard accretion governed by the MRI mechanism occurs alow for quantitative estimates of the parameter $\\omega$ and distinguishing of the KS naked singularities.  The most profound quantitative signatures of the KS naked singularities are related to the profiled spectral lines generated in both the outer and inner Keplerian discs. However, the signatures are strongly dependent on the inclination of the discs to the distant observer that should be known apriori, in order to obtain a convincing restrictions on the KS spacetime parameter. For objects close enough, the details of the appearance of the inner Keplerian discs could be observed giving direct qualitative signatures of the presence of KS naked singularity spacetimes.  Strong qualitative signatures of the KS naked singularities are expected in the case of the toroidal or spherical structures in the innermost parts of the near-extreme naked singularity spacetimes, being  influenced by the trapped photons. Of course, this is a theme for a future study."
    },
    "1402/1402.1766_arXiv.txt": {
        "abstract": "{We present four new epochs of $K\\sub{s}$-band images of the young pre-transitional disk around LkCa~15, and perform extensive forward modeling to derive the physical parameters of the disk.  We find indications of strongly anisotropic scattering (Henyey-Greenstein \\Rev{$g = 0.67^{+0.18}_{-0.11}$)} and a significantly tapered gap edge (`round wall'), but see no evidence that the inner disk, whose existence is predicted by the spectral energy distribution, shadows the outer regions of the disk visible in our images.  We marginally confirm the existence of an offset between the disk center and the star along the line of nodes; however, the magnitude of this offset (\\Rev{$x=27_{-20}^{+19}$}\\,mas) is notably lower than that found in our earlier $H$-band images \\citep{paper1}.  Intriguingly, we also find, at high significance, an offset of \\Rev{$y=69^{+49}_{-25}$}\\,mas perpendicular to the line of nodes.  If confirmed by future observations, this would imply a highly elliptical\\,---\\,or otherwise asymmetric\\,---\\,disk gap with an effective eccentricity of $e \\approx 0.3$.  Such asymmetry would most likely be the result of dynamical sculpting by one or more unseen planets in the system. Finally, we find that the bright arc of scattered light we see in direct imaging observations originates from the near side of the disk, and appears brighter than the far side because of strong forward scattering. } ",
        "introduction": "As potential indicators of planetary companions, transitional disks offer the tantalizing possibility of observing planet formation in action. They were initially identified as protoplanetary disks with a reduced near-infrared excess, indicating depleted inner regions \\citep{strom89,calvet05}. Millimeter-wave imaging has confirmed that these disks indeed contain large inner holes or annular gaps \\citep[e.g.,][]{andrews11}. However, optical and near-infrared images do not always reveal these holes \\citep{dong12}, indicating different spatial distributions of micrometer and millimeter-sized dust grains (see \\citealt{djo13}). In addition, stars continue to accrete substantial amounts of gas despite being cut off from the main gas reservoir in the outer disk \\citep{ingleby13,bergin04}, increasing the complexity of the puzzle.   To understand the complex geometry of these disks, multi-wavelength imaging is necessary. Only a handful of these gaps have been imaged in the optical/near-infrared \\Rev{\\citep[e.g.,][]{mayama12,hashimoto12,quanz13,avenhaus13,boccaletti13,garufi13}}, due to their high contrast ratio and proximity to the host star. A promising method to overcome these hurdles is angular differential imaging \\citep[ADI;][]{marois06}, which utilizes the natural field rotation of altitude-azimuth ground based telescopes to improve high-contrast imaging sensitivity. ADI efficiently reduces the impact of the stellar PSF wings, though at the cost of also imposing a non-negligible amount of self-subtraction for off-axis sources such as planets and disks. For point sources, this effect is easy to determine \\citep[e.g.,][]{lafreniere07,acorns}, but for extended sources such as disks, the subtraction effects are non-trivial and can affect the apparent morphology of the source. Forward modeling is a powerful tool for interpreting these images and extracting the disk geometry \\citep[e.g.,][]{thalmann11,milli12,thalmann13}.  One transitional disk that has been the focus of significant attention during the past few years is \\object{LkCa 15}, % a Sun-like host to a transitional disk with an inner hole the size of our Solar System \\citep[$\\sim$50~AU, e.g.,][]{pietu07}. Despite the large gap, LkCa 15 displays a significant near-infrared excess \\citep{espaillat07} and residual millimeter emission from small orbital radii \\citep{andrews11}, implying that an inner AU-sized, optically thick disk exists within the gap in addition to the outer disk \\citep{espaillat08,mulders10}. Sparse aperture masking observations have revealed an extended structure within the gap, which has been interpreted as a possible accreting protoplanet \\citep{kraus12}.  Recent simulations explore the structures induced into protoplanetary disks by individual embedded planets, and predict their observable signatures at near-infrared (NIR), mid-infrared, and sub-millimeter wavelengths \\Rev{\\citep{zhu11, dodson11, jangcondell13, djo13}}. \\Citet{kraus12} note that their planet candidate cannot be the body sculpting the wide gap, suggesting that additional planets may reside in the gap.  In 2010, we reported the first spatially resolved imaging of the LkCa~15 gap in the near infrared \\citep[][hereafter ``Paper I'']{paper1}. The observations confirmed the presence and size of the gap as being broadly consistent with that inferred by the \\Rev{spectral energy distribution (SED)} and seen at longer wavelengths, but also indicated a possible offset of the gap center from the location of the central star. This implied an eccentric gap edge, the likes of which have been suggested as a possible indication of planets within the gap, shepherding the disk material through their gravitational influence \\citep[e.g.,][]{kuchner03,quillen06}. The observations also revealed a strong brightness asymmetry between the Northern and Southern parts of the disk, which by itself could either be interpreted as preferential forward scattering from optically thin material at the near side of the disk, or reflection from an optically thick surface at the far side of the disk. The disk orientation derived by \\citet{pietu07} on the basis of asymmetries in their millimeter interferometry data favors the former scenario, whereas \\citet{jangcondell13} assume the latter scenario.  Overall, the very small angular scale of the disk and the prohibitive brightness contrast limited the amount of quantitative results that could be gleaned from the ADI images in \\citetalias{paper1} directly, necessitating the acquisition of deeper high-contrast imaging data as well as a comprehensive forward-modeling effort as a next step.  Here we present four new epochs of near-infrared high-contrast imaging of the \\Rev{LkCa~15} disk in the $K\\sub{s}$ band. Due to their superior Strehl ratio, the $K\\sub{s}$-band data offer cleaner disk images than the previously used $H$-band data.  Furthermore, the availability of several epochs of observation provides a more robust foundation for interpreting the disk morphology, through comparison of consistencies and scatter between the epochs. In order to extract quantitative results on the disk geometry from the imaging data, we generated an extensive parametric grid of model disks as described in \\citet{mulders10}, calculated their appearance in scattered light, forward-modeled them through the ADI process, and compare them to the data using a $\\chi^2$ metric.  In the following, we first describe our observations in detail, followed by a description of the data reduction procedure. We then discuss our modeling efforts to derive the properties of the disk and its inner hole, and subsequently show the results of this modeling. Finally, we discuss the implications of our results and present our conclusions. ",
        "conclusions": "Following up on our $H$-band results from \\citetalias{paper1} (epoch H1), we have obtained four new epochs of $K\\sub{s}$-band high-contrast direct-imaging observations of the LkCa~15 system in order to reveal new insights on the architecture of its pre-transitional disk. All four $K\\sub{s}$-band datasets (epochs K1--K4) were taken in consistent observing modes so as to enable direct comparison, using Gemini NIRI with pupil tracking to enable ADI data reduction.  In order to compensate for the known flux loss and morphological impact of ADI, we have performed an extensive forward-modeling analysis.  By simulating a parametric grid of scattered-light images of disk models, forward-modeling them through the ADI process, and comparing the results to the observed ADI images with a $\\chi^2$ metric, we have arrived at stringent quantitative constraints on the disk geometry.  These results are independent from, but consistent with, the constraints previously inferred from the system's SED \\citep{espaillat10} and from (sub-)millimeter interferometry \\citep{andrews11}.  However, they also include new findings that have only become accessible through near-infrared high-contrast imaging:  \\begin{itemize} \\item We find that the bright crescent seen in scattered-light ADI observations of LkCa~15 represents the strongly forward-scattering near side of the outer disk, rather than the far side (cf.\\ Section~\\ref{s:xo}).  \\item We tentatively confirm the existence of an asymmetry between the ansae as postulated in \\citetalias{paper1}.  Under the assumption that the disk can be described as azimuthally symmetric but physically displaced along the line of nodes (the major axis of the projected disk gap) by an offset $x$, we find consistently positive best-fit values for $x$ in all four $K\\sub{s}$-band epochs, in line with our H1 results (cf.\\ Section~\\ref{s:xo}). However, while the size of the best-fit $x$ offset is consistent with the H1 value in our lower-quality epochs K2 and K4, the higher-quality data of K1 and K3 indicate a much smaller offset consistent with zero.  This may suggest that the observed asymmetry may be caused by a different mechanism than a displaced disk gap, such as perhaps by a spiral density wave beyond the western ansa of the disk gap, which may be mistaken for the disk gap in lower-quality data.  \\item Furthermore, our best-fit models exhibit a surprisingly large disk offset $y$ perpendicular to the line of nodes (along the minor axis of the projected disk gap), bringing the brightly illuminated near-side rim of the outer disk closer to the star (cf.\\ Section~\\ref{s:y}).  Unlike $x$, the values of $y$ are consistent with each other and significantly inconsistent with zero in all four $K\\sub{s}$-band epochs, lending the result credibility.  While such a large displacement ($\\approx$0.3 times the gap radius, deprojected) constitutes a dramatic departure from the symmetric architecture commonly assumed for LkCa 15, it could plausibly have escaped detection in SED and submillimeter interferometry.  However, we note that a displaced circular gap is no longer a good approximation of an elliptical gap at an eccentricity of $e=0.3$, so the real shape of the LkCa~15 disk may well lie outside of the model parameter space considered in this analysis.  \\item We find evidence for a `round' gap wall; i.e.,\\ a gradual rather than abrupt gradient of surface density between the gap and the outer disk, consistent with the wall being sculpted by a planetary companion (cf.\\ Section~\\ref{s:w}).  \\item We obtain best results for a completely absent shadow from the inner disk on the outer disk wall (cf.\\ Section~\\ref{s:s}). This may indicate that the inner disk is inclined with respect to the outer disk, causing its shadow to miss the outer disk at most azimuths.  \\end{itemize}  \\noindent The best-fit disk models are visualized in Figure~\\ref{f:bestfit}, and the numerical results of the analysis are summarized in Table~\\ref{t:sanity}.  Finally, we acknowledge the fact that the results of our forward-modeling analysis are necessarily bound by the scope of the model and its parameter space, and that they are therefore simplistic approximations to the real, complex reality of the LkCa~15 disk. However, given the observational status quo, we consider them to represent the best estimate of the disk architecture to date.  In the near future, imaging polarimetry with next-generation adaptive optics facilities such as Subaru SCExAO \\citep{guyon11} or SPHERE ZIMPOL \\citep{thalmann08} may provide a more detailed view of the LkCa~15 disk with fewer model dependencies."
    },
    "1402/1402.4375_arXiv.txt": {
        "abstract": "Dark matter particles captured in the Sun would annihilate producing a neutrino flux that could be detected at the Earth. In some channels, however, the neutrino flux lies in the MeV range and is thus undetectable at IceCube, namely when the dark matter particles annihilate into $e^+e^-$, $\\mu^+\\mu^-$ or light quarks. On the other hand, the same interaction that mediates the annihilations into light fermions also leads, via higher order effects, to the production of weak gauge bosons (and in the case of quarks also gluons) that generate a high energy neutrino flux potentially observable at IceCube. We consider in this paper tree level annihilations into a fermion-antifermion pair with the associated emission of one gauge boson and one loop annihilations into two gauge bosons, and we calculate the limits on the scattering cross section of dark matter particles with protons in scenarios where the dark matter particle couples to electrons, muons or light quarks from the non-observation of an excess of neutrino events in the direction of the Sun. We find that the limits on the spin-dependent scattering cross section are, for some scenarios, stronger than the limits from direct detection experiments. ",
        "introduction": "\\label{sec:introduction}  The detection of a high energy neutrino signal from the Sun would strongly point towards exotic physics, most notably the annihilation of dark matter (DM) particles that have been captured in the solar interior. Conversely, the non-observation of a significant excess of events in neutrino telescopes in the direction of the Sun already sets fairly strong limits on the dark matter properties. For example, under the assumption that dark matter captures and annihilations are in equilibrium in the interior of the Sun, present measurements of the neutrino flux by the IceCube experiment set the 90\\% C.L. upper limit on the spin-dependent DM-proton interaction cross section $\\sigma^{\\rm SD}<1.3\\times 10^{-40}~(6.0\\times 10^{-39})~{\\rm cm}^2$ for a 250 GeV dark matter particle that annihilates into  $WW$ ($b\\bar b$) \\cite{Aartsen:2012kia}. Remarkably, for annihilation channels producing hard neutrinos (such as $WW$ or $\\tau\\tau$), the IceCube limits on the spin-dependent interaction cross section are more stringent than those reported by the most sensitive current experiments probing the same interaction, COUPP \\cite{Behnke:2012ys} and SIMPLE \\cite{Felizardo:2011uw}.  For some annihilation channels, however, the limits are much weaker, for instance when the dark matter annihilates into $e^+ e^-$, $\\mu^+\\mu^-$ or $q \\bar q$, with $q=u,d,s$. Due to the high density of matter in the interior of the Sun, where the annihilations take place, muons and light hadrons (such as pions and kaons) are quickly stopped before they decay. Therefore, the annihilation into muons or into light quarks produces neutrinos with an energy smaller than $\\sim 100$ MeV. Besides, annihilations into electrons do not produce neutrinos directly, although the interaction of the energetic electrons with the solar matter produces pions which  are in turn stopped before decaying, again producing low energetic neutrinos. All these annihilation channels then produce a neutrino flux with an energy well below the detection threshold of IceCube, although they could be detected at SuperKamiokande~\\cite{Bernal:2012qh,Rott:2012qb}.  On the other hand, in scenarios where the dark matter particle annihilates into a fermion-antifermion pair, final states with gauge bosons necessarily occur either from internal bremsstrahlung or from loop induced annihilations. Weak gauge bosons can decay producing energetic neutrinos while gluons produce heavy hadrons, which decay also producing energetic neutrinos, hence opening the possibility of observing in IceCube dark matter signals from the Sun in scenarios where the annihilation is driven by couplings to the electron, the muon or the light quarks. In this paper we will investigate this possibility, and derive limits on this class of scenarios from considering the high energy neutrino flux generated by the higher order processes of internal bremsstrahlung or loop annihilations.  While the annihilation cross sections for the higher order processes are expected to be small, under the common assumption that dark matter captures and annihilations are in equilibrium in the solar interior, the neutrino flux from the Sun is determined, not by the annihilation cross sections, but by the capture rate and by the branching fractions of the different annihilation channels. We will argue that, for some scenarios, the branching fraction for the higher order annihilation processes with gauge bosons in the final state is sufficiently large to produce a high energy neutrino flux at the reach of IceCube, for values of the spin-dependent interaction cross section in agreement with the limits set by the direct search experiments COUPP and SIMPLE.  The paper is organized as follows. In Section \\ref{sec:Contact_Interactions} we undertake a model independent analysis where the tree level dark matter annihilation into a fermion-antifermion is mediated by a contact interaction, and consider separately the limits on the dark matter interaction cross section with protons in scenarios where the most important source of high energy neutrinos is either the annihilation into a fermion-antifermion pair with the associated emission of a gauge boson or the one loop annihilation into two gauge bosons. In Section \\ref{sec:toy-model} we focus on a toy model with a Majorana dark matter particle that couples to a light fermion via a Yukawa coupling with a scalar, and where the branching ratios for the two-to-three and the one loop annihilations are calculable in terms of the parameters of the model. Lastly, in Section \\ref{sec:Conclusions} we present our conclusions. We also include appendices including fitting formulas for the (anti-)neutrino spectra generated by the final state radiation of a gauge boson, tables with the upper limits on the scattering cross section assuming annihilations into $ZZ$ and $gg$ for various dark matter masses, and the expressions for the annihilation cross sections in the relevant channels in the toy model discussed in Section \\ref{sec:toy-model}. ",
        "conclusions": "\\label{sec:Conclusions}  We have investigated the impact of higher order annihilation processes in the generation of a high energy neutrino flux from dark matter annihilations in the Sun in scenarios where the dark matter particle couples to the electron, the muon or to the light quark. We have argued that, while the annihilation into a fermion-antifermion pair generates only MeV neutrinos, the associated emission of a gauge boson off the final state and the loop induced annihilation into two gauge bosons do generate a high energy neutrino flux, thus opening the possibility of probing these scenarios at IceCube.  We have first adopted a model independent approach where the annihilation into the fermion-antifermion is induced by a contact interaction and we have considered two limiting scenarios depending on which is the most important source of high energy neutrinos, whether the loop annihilations or the final state radiation of gauge bosons. In both cases, we have derived limits on the spin-dependent and spin-independent dark matter interaction cross sections with protons under the assumption that captures and annihilations are in equilibrium in the solar interior. We have found fairly stringent limits for the spin-dependent scattering cross section for dark matter masses between 100 GeV and 10 TeV which are complementary to the limits from the direct search experiments COUPP and SIMPLE and, in the case of coupling to leptons, stronger.  Lastly, we have carefully analyzed the neutrino signals in a toy model consisting in a Majorana fermion as dark matter particle that couples to a right-handed electron or first generation quark  via a Yukawa coupling with a scalar. The branching ratios of the different processes are calculable in this model: the two-to-two annihilation into a fermion-antifermion pair is helicity- and velocity-suppressed, and always negligible, while the two-to-three and loop induced annihilations into weak gauge bosons have a sizable branching fraction, their relative importance being dependent on the parameters of the model. More specifically, the two-to-three process dominates when the scalar that mediates the interaction is degenerate in mass with the dark matter particle, while the loop process dominates when there is a large mass hierarchy between them. We have found that the limits on the interaction cross section depend mildly, at most by a factor of two, on the mass of the scalar particle that mediates the interaction. Furthermore, we have found stringent limits on the spin-dependent interaction cross section for the model where the dark matter particle couples to right-handed electrons, which are stronger than the limits set by direct search experiments.  \\vspace{0.5cm}"
    },
    "1402/1402.6222_arXiv.txt": {
        "abstract": "In this paper, we describe a hybrid-extended Kalman filter algorithm to synchronize the clocks and to precisely determine the inter-spacecraft distances for space-based gravitational wave detectors, such as (e)LISA. According to the simulation, the algorithm has significantly improved the ranging accuracy and synchronized the clocks, making the phase-meter raw measurements qualified for time-delay interferometry algorithms. ",
        "introduction": "The Laser Interferometer Space Antenna (LISA) \\cite{LISA98,Danzmann03,LISA11} is a space-borne gravitational wave (GW) detector, aimed at various kinds of GW signals in the low-frequency band between $0.1\\,$mHz and $1\\,$Hz. It consists of three identical spacecraft (S/C), each individually following a slightly elliptical orbit around the Sun, trailing the Earth by about $20^\\circ$. These orbits are chosen such that the three S/C retain an equilateral triangular configuration with an arm length of about $5\\times10^9\\,$m as much as possible. This is accomplished by tilting the plane of the triangle by about $60^\\circ$ out of the ecliptic. Graphically, the triangular configuration makes a cartwheel motion around the Sun. As a (evolving) variation of LISA, eLISA \\cite{Danzmann13} is an ESA L2/L3 candidate space-based GW detector. It consists of one mother S/C and two daughter S/C, separated from each other by $1\\times10^9\\,$m. Although the configurations are slightly different, the principles and the techniques are equally applicable. Therefore, we will mainly focus on LISA hereafter.  Since GWs are propagating spacetime perturbations, they induce proper distance variations between test masses (TMs) \\cite{Carbone}, which are free-falling references inside the S/C shield. LISA measures GW signals by monitoring distance changes between S/C. Spacetime is very stiff. Usually, even a fairly strong GW still produces spacetime perturbations only of order about $10^{-21}$ in dimensionless strain. This strain amplitude can introduce distance changes at the pm level in a $5\\times10^9\\,$m arm length. Therefore, a capable GW detector must be able to monitor distance changes with this accuracy. The extremely precise measurements are supposed to be achieved by large laser interferometers. A schematic classic LISA configuration with exchanged laser beams is shown in Fig.~\\ref{fig:LISA_const}. LISA makes use of heterodyne interferometers with coherent offset-phase locked transponders \\cite{McNamara05}. The phasemeter \\cite{Shaddock06} measurements at each end are combined in postprocessing to form the equivalent of one or more Michelson interferometers. Information of proper-distance variations between TMs is contained in the phasemeter measurements.  Unlike the several existing ground-based interferometric GW detectors~\\cite{LIGO92,GEO02,VIRGO97}, the armlengths of LISA are varying significantly with time due to celestial mechanics in the solar system. As a result, the arm lengths differ by about $1\\%$ ($5\\times10^7\\,$m), and the dominating laser frequency noise will not cancel out. The remaining laser frequency noise would be stronger than other noises by many orders of magnitude. Fortunately, the coupling between distance variations and the laser frequency noise is very well known and understood. Therefore, we can use time-delay interferometry (TDI) techniques~\\cite{TDIexperiment,TDI03,TDI99,TDI03b,TDI04,TDI05,TDI02,TDI05b,TDI12}, which combine the measurement data series with appropriate time delays, in order to cancel the laser frequency noise to the desired level.  However, the performance of TDI \\cite{Petiteau08,TDI05} depends largely on the knowledge of armlengths and relative longitudinal velocities between the S/C, which are required to determine the correct delays to be adopted in the TDI combinations. In addition, the raw data are referred to the individual spacecraft clocks, which are not physically synchronized but independently drifting and jittering. This timing mismatch would degrade the performance of TDI variables. Therefore, they need to be referred to a virtual common ``constellation clock\" which needs to be synthesized from the inter-spacecraft measurements. Simultaneously, one also needs to extract the inter-spacecraft separations and synchronize the time-stamps properly to ensure the TDI performance. More precisely, the knowledge of the distances between S/C needs to be better than 1 m rms at 3 Hz. Accordingly, the differential clock errors between the S/C are required to be estimated to a precision of 3.3 ns rms at 3 Hz\\footnote{ Better knowledge of the armlengths and the differential clock errors will result in better cancellation of laser frequency noise in TDI variables, since the residual laser frequency noise in TDI variables is proportional to the armlength errors and the differential clock errors.}.  These are the main goal of the first stage of LISA data processing, which is the main topic of this paper.  The paper is organized as follows. In the next section, we will describe the entire LISA data processing chain, identifying the first stage of LISA data analysis. In Sec.~\\ref{sec:Meas1} and Sec.~\\ref{sec:Meas2}, we will introduce and formulate the inter-spacecraft measurements. In Sec.~\\ref{sec:Kalman} and \\ref{sec:LISAmodel}, we describe the hybrid-extended Kalman filter algorithm and design a Kalman filter model for LISA. In Sec.~\\ref{sec:Simulation}, we show the simulation results. Finally the summary comes in Sec.~\\ref{sec:Sum}.  \\begin{figure} \\includegraphics[width=0.5\\textwidth]{constellation.pdf} \\caption{ \\label{fig:LISA_const} Schematic configuration of LISA S/C and the exchanged laser beams (by S. Barke \\cite{LISAreport}).} \\end{figure} ",
        "conclusions": "\\label{sec:Sum}   We have modeled LISA inter-spacecraft measurements and designed a hybrid-extended Kalman filter to process the raw measurement data.  In the designed Kalman filter model, there are 24 variables in the state vector and 18 variables in the measurement vector. Therefore, (i) the state vector, in principle, cannot be fully determined from the measurements, which is one of the major differences from the global positioning system (GPS)~\\cite{GPS,Mohinder08} tracking problem, where the number of measurements is larger than the number of unknowns in the state vector. Other important differences from GPS are: (ii) the position and the time of the emitter S/C are unknown (more specifically, they are also to be estimated). Therefore, the measurements associated with each laser link are functions of the receiver S/C at the current time and the emitter S/C at a previous time. This gives rise to the ``causality\" problem that the inference of current S/C depends also on future measurements. (iii) The measurements recorded on different S/C are unsynchronized, and contaminated by different unknown clock jitter. These differences make the first stage of LISA data processing much more challenging.  This paper presents a major step towards a high fidelity end-to-end simulation of the entire LISA data processing chain. We have identified the problems, established the framework and formulated inter-spacecraft measurements that are crucial to the first stage data processing. Two important effects have not yet been included in the current simulation: (i) the time delay is only partly simulated. In the simulation, the ranging measurements consist of the time delay and other noise. However, the dependence of inter-spacecraft measurements on the emitter S/C at a delayed time is simulated as that at the current time. Therefore, the effects of the `causality' problem and the Sagnac differential delay do not present in this simulation. These issues are being investigated in our follow-on work \\cite{Wang14b}, where we find out that it is more appropriate to simulate these effects in the full-relativistic framework. (ii) The clock jitter is not included in the recording time yet, but only in the measurements. This effect is included in our follow-on work \\cite{Wang14b}.  The current simulation shows that the hybrid-extended Kalman filter can well decouple the arm lengths from the clock biases and significantly improve the relative measurements, such as arm lengths, relative clock jitters and relative frequency jitters etc. However, the absolute variables in the state vector cannot be determined accurately. These variables include the absolute positions and velocities of the spacecraft, the absolute clock drifts and the absolute frequency drifts. This is mainly due to the fact that only the differences are measured and the number of measurements is lower than the number of variables in the state vector.  It can be better understood by taking a closer look at the measurement equations~\\ref{eq:R}, \\ref{eq:D} and \\ref{eq:C}. In fact, only the relative positions $\\sqrt{(x_j-x_i)^2+(y_j-y_i)^2+(z_j-z_i)^2}$ and the relative longitudinal velocities $(\\vec{v}_j-\\vec{v}_i)\\cdot \\hat{n}_{ij}$ appear in the measurements. Neither absolute positions nor absolute velocities are directly measured. Thus, it is impossible to fully constrain the entire LISA configuration only with these inter-spacecraft measurements. The clock jitters only appear in Eq.~\\ref{eq:R} in the form of $\\delta T_j - \\delta T_i$, which means the common clock drifts are undetermined. The relative USO frequency jitters $\\delta f_j - \\delta f_i$ are measured in Eq.~\\ref{eq:C}. The absolute USO frequency jitters $\\delta f_j$ appear in Eq.~\\ref{eq:D}. However, $\\delta f_j / f_j^\\textrm{nom}$ is far less than 1, hence Eq.~\\ref{eq:D} can provide only very limited information about $\\delta f_j$. As a result, the absolute USO frequency jitters $\\delta f_j$ are poorly determined.       \\appendix"
    },
    "1402/1402.6014_arXiv.txt": {
        "abstract": "We show that precision measurement of (1) sum of neutrino masses by cosmological observation and (2) lifetime of neutrinoless double beta decay in ton-scale experiments, with supplementary use of (3) effective mass measured in single beta decay experiment, would allow us to obtain information on the Majorana phase of neutrinos. To quantify the sensitivity to the phase we use the CP exclusion fraction, a fraction of the CP phase parameter space that can be excluded for a given set of assumed input parameters, a global measure for CP violation. We illustrate the sensitivity under varying assumptions, from modest to optimistic ones, on experimental errors and theoretical uncertainty of nuclear matrix elements. Assuming that the latter can be reduced to a factor of $\\simeq 1.5$ we find that one of the two Majorana phases (denoted as $\\alpha_{21}$) can be constrained by excluding $\\simeq 10-40\\%$ of the phase space at $2\\sigma$ CL  even with the modest choice of experimental error for the lowest neutrino mass of 0.1 eV. The characteristic features of the sensitivity to $\\alpha_{21}$, such as dependences on the true values of $\\alpha_{21}$, are addressed. ",
        "introduction": "\\label{sec:intro}  After all the angles in the lepton mixing matrix \\cite{MNS} and the mass squared differences of neutrinos are determined by the atmospheric, the solar, the reactor, as well as the accelerator neutrino oscillation experiments, we have entered into a new phase of neutrino physics~\\cite{Kajita:2012vc,McDonald:2004dd,Inoue:2004wv}. Of course, we are still with the important unknowns in the lepton sector, CP violating phase of the Kobayashi-Maskawa type~\\cite{KM} and the neutrino mass ordering. Yet, we do have relatively clearer view of how these questions are to be settled, either through the ongoing searches, or by new measurement to be carried out by the powerful apparatus described in various experimental proposals to date.  It is quite possible that even after these unknown quantities are measured by experiments the questions of somewhat different category will still remain: What are the absolute masses of neutrinos? Are neutrinos Dirac or Majorana particles? If Majorana, which values of the Majorana CP phases \\cite{Schechter:1980gr,Bilenky:1980cx,Doi:1980yb} are chosen by nature? The first one, knowing the absolute mass scale, is crucial to understand physics behind the origin of neutrino mass, and the latter two are likely the key to understand baryon number asymmetry in the universe \\cite{leptogenesis}.  In this paper, we focus on detectability of the Majorana CP phase in the light of informations to be obtained by cosmological observation and neutrino experiments in the coming precision era. In particular, in the near future the double beta decay experiments aim to cover entire region of parameters allowed by the inverted mass ordering,\\footnote{ Throughout the paper the ``ordering'' of masses refers to that of the ones relevant for atmospheric neutrino oscillation, namely, normal (inverted) mass ordering corresponds to $m_2 < m_3$ ($m_2 > m_3$). } see Fig.~\\ref{fig:m0nbb-m0}. We show that given the future perspective, when combined with precision measurement in cosmology, one can indeed expect some sensitivities, though limited ones, to one of the phases (denoted as $\\alpha_{21}$, see section~\\ref{subsec:double-beta}). We hope that this result can trigger renewed interests in the problem of detectability of the Majorana phase despite the dominant pessimism which prevailed in the past decades. See, e.g., Refs.~\\cite{Minakata:1996vs,Bilenky:2001rz,Czakon:2001uh,Pascoli:2001by,Barger:2002vy,Nunokawa:2002iv,Pascoli:2002qm,Deppisch:2004kn,Joniec:2004mx,Pascoli:2005zb,Choubey:2005rq,Simkovic:2012hq} for the foregoing efforts which examined the sensitivity to the Majorana phase.  We feel that the present moment is the particularly right time to revisit these issues because of the following reasons. Firstly, cosmological observations entered into the precision era so that their data became sensitive to the absolute masses of neutrinos better than the current bounds from laboratory experiments, as seen most notably by the results obtained by the Planck collaboration~\\cite{Ade:2013zuv}. See also Ref.~\\cite{Abazajian:2011dt} for overview and many relevant references. It makes us possible to employ a new strategy that the absolute neutrino mass scale can be constrained mainly by cosmology and we can use data of neutrinoless double beta decay experiments, with supplementary use of single beta decay data, to constrain the Majorana phases of neutrinos. Even more interestingly a discrepancy between the cosmic microwave background (CMB) and lensing observations about cluster correlations would prefer massive neutrinos with mass of a few tenth of eV \\cite{Battye:2013xqa}. If confirmed, it would make the basis of our analysis more robust.  Secondly, as mentioned above, several new neutrinoless double beta decay experiments are in line to measure or constrain, the ``ee'' element of neutrino mass matrix in the flavor basis, so called the effective Majorana mass, to a few tens of meV~\\cite{Albert:2014awa,Auger:2012ar,Gando:2012zm,Agostini:2013mzu,Arnaboldi:2002du,Hartnell:2012qd,Gaitskell:2003zr,Arnold:2010tu,Alvarez:2011my}. For reviews and description of the other projects, see e.g., Refs.~\\cite{GomezCadenas:2011it,Giuliani:2012zu,Cremonesi:2013bma}. Needless to say, they are the prime candidates among various experiments that can measure or constrain the Majorana phases. We should also note that intensive efforts have been devoted in order to improve the calculations of the nuclear matrix element (NME) which is crucial to determine the value of effective Majorana mass from the measured half life time in double beta decay experiments~\\cite{Rodin:2003eb,Rodin:2007fz,Caurier:1996zz,Caurier:2007wq,Simkovic:2009pp,Kortelainen:2007rh,Kortelainen:2007mn,Barea:2009zza,Barea:2012zz,Barea:2013bz,Menendez:2008jp}. See, also Ref.~\\cite{Cremonesi:2013bma} for review of the recent progress in the calculations of NME.  Thirdly, the uncertainties in the relevant lepton mixing angles get decreased dramatically recently. The error of $\\sin^2 \\theta_{12}$ is now only $\\sim 4\\%$\\cite{Gando:2013nba}, and moreover it will be decreased to a sub-percent level by the future medium-baseline reactor experiments~\\cite{Kettell:2013eos,Seo:2013swa}. The mixing angle $\\theta_{13}$ whose value was unknown until recently before intriguing indication from T2K~\\cite{T2K} by $\\nu_{e}$ appearance events, is now precisely measured \\cite{An:2012bu,RENO}. The error of $\\sin^2 \\theta_{13}$ will soon reach to a even greater precision, $\\sim 5\\%$ level. See Fig.~\\ref{fig:m0nbb-m0} which shows how small is the effect of the current uncertainties of the mixing parameters onto the allowed regions by double beta decay experiments. These developments are important to tighten up the constraint on Majorana CP phases and to place it on a firmer ground.  In addition, the single beta decay experiment KATRIN \\cite{Wolf:2008hf} (see also \\cite{KATRIN}) will constrain neutrino masses to a level of 0.2 eV at 90 \\% CL, or observe the effects consistent with masses of this order, in a manner quite complementary to cosmological observations. While enjoying the above new inputs, our analysis is nothing but a continuation of the numerous similar analyses attempted by the many authors, from which we are certainly benefited.  To display the sensitivities to the Majorana CP phases in a global way, we use in this paper the CP exclusion fraction $f_{\\rm CPX}$ \\cite{Machado:2013kya}, which is defined as a fraction of the CP phase space that can be excluded at a certain CL for a given set of input parameters. We believe that it is a useful tool particularly in the initial era of search for effects of the Majorana CP phase, in which even the partial exclusion of allowed range of the phase would be very valuable.\\footnote{ On the other hand, $f_{\\rm CPX}$ has intimate connection to the error of CP phase, an appropriate measure in precision measurement era, as discussed in \\cite{Machado:2013kya}. It suggests that the CP exclusion fraction can be a universally usable measure for CP sensitivity. } In \\cite{Machado:2013kya} $f_{\\rm CPX}$ was used to display sensitivity to the lepton Kobayashi-Maskawa phase achievable in the ongoing long-baseline neutrino oscillation experiments. For earlier discussion of a CP sensitivity measure closely related to $f_{\\rm CPX}$ see \\cite{Winter:2003ye,Huber:2004gg}.  In the following two sections, \\ref{sec:observables} and \\ref{sec:analytic-estimation}, we briefly review the basic features of the observables used and describe simple analytic estimation of the effects of Majorana phases. The expert readers must go directly to the analysis results given in sections~\\ref{sec:why} and \\ref{sec:results-exclusion-fraction}, starting from describing the analysis method in section~\\ref{sec:analysis-method}. ",
        "conclusions": "\\label{sec:conclusions}  Assuming the mass mechanism of $0\\nu\\beta\\beta$ decay, we have studied in this paper how and to what extent the Majorana phases can be constrained in the light of future accurate measurements of $0\\nu\\beta\\beta$ decay rate and an independent precision measurement of neutrino's absolute mass scale. In our estimates of the errors described in Sec.~\\ref{sec:exp-errors} cosmological observation of the sum of neutrino masses $\\Sigma$ plays the latter role. We have demonstrated that even in the situation that double beta decay experiment alone cannot say anything about the Majorana phase, it is possible to obtain highly non-trivial constraints on the Majorana phase, if accurate measurement of $\\Sigma$ is added. Indeed, it is unlikely that $0\\nu\\beta\\beta$ decay experiment by itself can place any useful constraints on the Majorana phase within our assumption of the experimental errors. Notice, however, that it is not because they suffer from the NME uncertainty (which of course makes the situation worse), but because they lack independent key information of the absolute neutrino mass scale in itself (see Sec.~\\ref{how-constrained}). The remarkable feature of our results is that the synergy between cosmology and double beta decay experiments in constraining the Majorana phase is quite visible.  The general tendency of the sensitivity to the Majorana phases are as follows: \\begin{itemize}  \\item Dependence of $m_{0\\nu\\beta\\beta}$ on $\\alpha_{31}$, one of the two Majorana phases, is very weak due to the suppression by $s^2_{13}$. Hence, little sensitivity can be expected for $\\alpha_{31}$. The sensitivity to $\\alpha_{21}$ depend on $\\alpha_{31}$, but only very weakly.  \\item The regions of the best and the worst sensitivities to $\\alpha_{21}$ are located at around the true values $\\alpha_{21} \\simeq 0$ or $\\pi$, and at $\\alpha_{21} \\simeq {2\\pi}/{3}$, respectively, in agreement with our analytic estimate in Sec.~\\ref{subsec:degenerate}.  \\item For both mass orderings and within any uncertainties of the NME used, the better sensitivity to the CP phase $\\alpha_{21}$ is obtained for larger value of $m_0$, apart from minor exception of region $m_0 \\lsim 0.05$ eV which is mostly in the horizontal branch in Fig.~\\ref{fig:m0nbb-m0} for the inverted mass ordering.  \\end{itemize}  We took as the reference set up of the uncertainties that the effective masses observed in $0\\nu\\beta\\beta$ and $\\beta$ decay experiments can be measured with the errors of 0.01 eV and 0.06 eV, respectively, whereas the sum of neutrino masses $\\Sigma$ can be determined with the uncertainty of 0.05 eV by cosmological observation. We have assumed the uncertainty factor $r_{\\text{\\tiny NME}}$ (defined as the ratio of the maximum and minimum values of the theoretically expected NME values) of a range $1.3 - 2.0$. We have treated it in a different way from these experimental errors, in similar way as the treatment of flux normalization uncertainty. To show the dependence of these assumed errors on sensitivity to Majorana phase, we have also examined the case with more optimistic errors, 0.02 eV for the error of $\\Sigma$ and $r_{\\text{\\tiny NME}} =1.1$ for the NME uncertainty.  To display the sensitivities to the Majorana phase globally we have used the new measures, the CP exclusion fraction, $f_{\\rm CPX}$, a fraction of the CP phase space that can be excluded for a given set of input parameters. $f_{\\rm CPX}$ is a useful global measure, as the CP fraction widely used in the sensitivity studies for the long-baseline neutrino oscillation experiments is, but in a different way. We believe that $f_{\\rm CPX}$ is better suited to reveal (relatively poor) sensitivities in the early era of the search for the CP phase effect, as discussed in \\cite{Machado:2013kya}. Since measuring the Majorana phase is very challenging in any means, we believe that even the partial exclusion of its possible range is quite useful.  We have shown with our reference setup and $r_{\\text{\\tiny NME}} =1.5$ that $\\alpha_{21}$ can be constrained by excluding $\\simeq 10-40\\%$ of the phase space of $\\alpha_{21}$ at $2\\sigma$ CL for the lowest neutrino mass of 0.1 eV for both mass orderings. Even if we use 0.02 eV for the error of $\\Sigma$ the excluding fraction does not change much, and the exclusion fraction is $\\simeq 10-50\\%$. Fortunately, the $f_{\\rm CPX}$ contours are rather stiff against change in confidence level from $2\\sigma$ to $3\\sigma$, indicating robust nature of sensitivity to the Majorana phase. We also should note that the sensitivity to $\\alpha_{21}$ becomes significantly better when the uncertainty of the nuclear matrix element ($r_{\\text{\\tiny NME}}$) is reduced from the factor 2 to 1.5, independent of the assumed mass ordering and value of $\\alpha_{31}$.  It is only recently that such a discussion started to obtain real perspective because cosmology entered into the precision era, as demonstrated most notably by an epoch making precision measurement achieved by Planck. Yet, the accuracy of measurement of the sum of neutrino masses which warrants reasonably strong restriction to the Majorana phase is rather demanding, $\\sigma_{\\Sigma} \\simeq 0.02$ eV, which probably requires the Euclid satellite as well as the next generation galaxy surveys, in addition to the better understanding of the different systematics exist in the different set of the cosmological data. Of course, it also requires accurate measurement of lifetime of $0\\nu\\beta\\beta$ decay in a ton scale experiment with very low background, which would allow uncertainty in measurement of $m_{0\\nu\\beta\\beta}$ as small as $\\simeq 0.01$ eV.  One may argue that too much relying on cosmological observation in determining $\\Sigma$ is dangerous because, neutrinos being a minor player in the universe, its precise determination is possible only in a model dependent way. Though it is a valid point we must bear in mind that future precision observation itself offers an even more stringent test of the $\\Lambda$CDM paradigm. Then, it may be conceivable that the SM of cosmology could be eventually born out from such process, which we assumed as a prerequisite of our analysis. It would allow us a reliable measurement of the neutrino component in the universe and to determine the sum of their masses with a sufficient robustness. The discussion in this paper would serve as a prototype of the similar analysis that should be done after we have acquired the established standard cosmological model.  \\appendix"
    },
    "1402/1402.5151_arXiv.txt": {
        "abstract": "We present the first systematic comparison of ultraviolet-millimeter spectral energy distributions (SEDs) of observed and simulated interacting galaxies. Our sample is drawn from the Spitzer Interacting Galaxy Survey, and probes a range of galaxy interaction parameters. We use 31 galaxies in 14 systems which have been observed with \\emph{Herschel}, \\emph{Spitzer}, \\emph{GALEX}, and 2MASS. We create a suite of \\gadgetthree hydrodynamic simulations of isolated and interacting galaxies with stellar masses comparable to those in our sample of interacting galaxies. Photometry for the simulated systems is then calculated with the \\sunrise radiative transfer code for comparison with the observed systems.  For most of the observed systems, one or more of the simulated SEDs match reasonably well. The best matches recover the infrared luminosity and the star formation rate of the observed systems, and the more massive systems preferentially match SEDs from simulations of more massive galaxies. The most morphologically distorted systems in our sample are best matched to simulated SEDs close to coalescence, while less evolved systems match well with SEDs over a wide range of interaction stages, suggesting that an SED alone is insufficient to identify interaction stage except during the most active phases in strongly interacting systems. This result is supported by our finding that the SEDs calculated for simulated systems vary little over the interaction sequence. ",
        "introduction": "Galaxy interactions, especially major mergers, are responsible for some of the most dramatic activity seen in galaxies. In the canonical view, interactions stimulate star formation, thereby powering the high infrared (IR) luminosities often seen in such systems \\citep[e.g.,][]{vei02}: driving gas inflows to the central regions, resulting in heightened activity of the central supermassive black hole and local starburst activity \\citep[e.g.,][]{dim05, spr05b}, and leading to significant morphological distortions \\citep[e.g.,][]{hop06, mih94, mih96}. These activities, however, occur over timescales that make detecting evolution in individual systems or tracing the corresponding development in physical processes impossible. Hydrodynamic simulations of interacting galaxies provide a means of probing the interaction sequence and bypassing the problem of the timescales.  A crucial test of any simulation is its ability to reproduce observations. Hydrodynamic simulations of galaxy interactions have primarily been tested in two ways:  by how well they reproduce the (optical) morphological distortions seen in such systems, and by how closely their simulated emission matches that of real systems. Some simulations are designed to reproduce specific systems \\citep[e.g.,][]{privon13, karl13}, while others compare specific properties, such as colors, of a suite of simulations to observations \\citep[e.g.,][]{sny13, jon10}.  \\citet{too72} were the first to systematically model the morphologies of interacting galaxies. They used simple simulations of massless particles around two masses  to reproduce the tails and bridges seen in systems like the M51, the Mice (NGC\\,4676) and the Antennae (NGC\\,4038/4039). Much more recently, \\citet[see also Barnes 2011]{bar09} developed Identikit, a modeling tool that uses N-body simulations to reproduce the morphology and kinematics of tidal tails in interacting systems. \\citet{privon13} demonstrated Identikit's ability to reproduce the morphology and H I kinematics of NGC\\,5257/5258, the Mice, the Antennae, and NGC\\,2623 and to estimate the time since the first pericenter passage and until coalescence. \\nocite{bar11}  Morphological analyses like these inherently suffer from an obvious bias: simulations trace mass but observations trace light. Better comparisons propagate light from the simulated luminous matter to a fiducial observer. \\sunrise  \\citep{jon06} accomplishes exactly that. It is a radiative transfer code that propagates the emission of simulated stars and active galactic nuclei (AGN) through a dusty interstellar medium (ISM) generated by the hydrodynamic simulations. It is an ideal tool for creating simulated spectral energy distributions (SEDs) for comparison to photometry. For example, \\citet{jon10} simulated the SEDs of seven isolated galaxies, which they compared to Spitzer Infrared Nearby Galaxies Survey \\citep[SINGS;][]{ken03} galaxies from \\citet{dale07}. The \\citet{jon10} simulations did not cover all of the parameter space spanned by SINGS; nonetheless, good matches from the SINGS sample were found for each of the simulated galaxies, demonstrating the ability of \\sunrise to produce realistic galaxy SEDs. \\citet{karl13} combined analyses of the morphology and emission, by creating a set of hydrodynamic simulations to reproduce the morphology of the Antennae and performing radiative transfer to determine the predicted emission in the Herschel Space Observatory's Photodetector Array Camera and Spectrometer (PACS) bands.  This paper is the first systematic comparison of the observed and simulated SEDs spanning the range from the ultraviolet (UV) to the far-IR (FIR) for interacting galaxies. \\citet{jon10} determined which of the SINGS galaxies were best reproduced by the SEDs of their simulated isolated spiral galaxies.  Our study takes a related but different approach: we do not create simulations aimed at specifically reproducing our observed interactions; rather, we produce a range of simulated interactions and examine how realistically they reproduce observed systems, determine which of the  simulated SEDs best reproduce the observed SEDs of interacting systems, and identify the simulation properties, such as stellar mass, SFR, or interaction stage, common to the set.  At high redshifts, morphological details become impossible to resolve and so estimates of interaction stage based on morphology  \\citep[e.g.,][]{dopita02} become unworkable. A spectral marker for interaction stage would be a powerful tool for examining how  interactions at high redshift compare to local interactions. Therefore, we ask whether there is an unambiguous signature of the interaction stage in the SED. In this paper, we compare the SEDs of a suite of simulations of interacting and isolated spiral galaxies to the SEDs of  31 interacting galaxies to examine the simulation properties able to reproduce the SED of an observed system. A clear extension of our study is to test whether there is a signature of the morphology in the SED by finding common morphology either within the set of best matches or between the matches and the observation.  This article is organized as follows. We summarize our sample selection and the photometry  in Section 2. In Section 3, we describe the hydrodynamic simulations and the radiative transfer done in post-processing.  We discuss our matching methodology and the best-matched SEDs in Section 4. Section 5 contains a  discussion of the origins of the best- and worst-matched SEDs,  a comparison between the stellar and dust masses,  dust luminosity, star formation rate (SFR), specific star formation rate (sSFR) of the observed systems and the best-matched simulated counterparts, an analysis of the effectiveness of morphology-based interaction stage classification scheme, and an examination of the evolution of SEDs in major mergers. We summarize our results in Section 6.  \\begin{figure}[bhp] \\centerline{\\includegraphics[width=\\linewidth]{f1_color.pdf}} \\caption{The range of IR luminosity and stellar mass covered by the observed sample used in this paper (red diamonds), the full parent SIGS sample (black stars), and the simulations (boxes in blue (interactions) and black (isolated galaxies)). The simulations cover the full range of the observed IR luminosity, and their mass range spans more than two orders of magnitude and are representative of most of the observed sample. The widths and height of a simulation box shows the range of the IR luminosity and stellar mass over the simulation. } \\label{par_range} \\end{figure} ",
        "conclusions": "We presented the first systematic comparison of SEDs of observed and simulated interacting galaxies. Our sample of 31 galaxies was observed in up to 25 bands including \\emph{GALEX}, \\emph{Spitzer}, and \\emph{Herschel}. We created a suite of \\gadgetthree hydrodynamic simulation of four galaxies evolving in isolation and the ten pair-interactions evolved from first passage through coalescence to the post-merger stage. Simulated SEDs were calculated using \\sunrise for two different treatments of the ISM's multiple phases, and sets of the best-matching of these SEDs were determined for each observed pair. From our comparison of the simulated and observed SEDs, we conclude that:\\\\  (1)~~For most observed systems, at least one mock SED from the simulations provides a reasonably good (often statistically acceptable) fit to the observed SED. The best-matched SEDs generally come from the same simulation, often the major mergers. They tend to cluster around coalescence and the constraint in timing is tightest for the most evolved systems. We did not find, however, that best matches preferentially come from a certain viewing angle.  (2)~~Neither treatment of the sub-resolution ISM is preferred for our sample as a whole. Some interactions are better matched when only the diffuse dust was used in the radiative transfer calculations (DISM) while others have SEDs which are better reproduced especially in the FIR when the total dust mass is used in the calculations (AISM). Half of the observed interactions are equally well matched by the two sets of simulated SEDs.  (3)~~The best matches recover IR luminosity and SFR fairly well.  (4)~~Stellar and dust masses show indications that more massive (or dustier) systems tended to be matched by simulations of more massive (or dustier) galaxies, but greater coverage of the simulated parameter space is necessary to reliability recover stellar and dust masses.  (5)~~The DISM SEDs have dust that is typically hotter than the observed systems. The temperature is better recovered with the AISM matches.  (6)~~Our SED matching techniques is able to reliably identify the interaction stage of the strongly interacting systems near coalescence. In contrast, the less strongly interacting morphology-based classes cover a wide range of interaction stages in their best matches. This suggests that the integrated SED alone is typically insufficient to identify the interaction stage, except for the most strongly interacting systems and the passively evolving merger remnants.  (7)~~The SEDs of the simulated systems in the different stages are generally quite similar, supporting our previous conclusion. The two exceptions are first, the passively evolving merger remnants (which exhibit markedly less UV and FIR emission because of quenching), and second, the most strongly interacting systems. For the latter, the SED varies significantly because of the rapid variation in the SFR and AGN luminosity. \\\\  Our study improved upon the comparison of \\citet{jon10} who tested  \\sunrise against the SEDs of the SINGS galaxies by addressing two of the five issues of concern they raised: the small range of simulations available to them to fit normal galaxies, and the lack of treatment for the cold ISM. Regarding their first issue, we found in our larger range of systems that the simulations and \\sunrise do a good (albeit not perfect) job of fitting observed SEDs over a much larger range than observed with SINGS. Additionally, our study focused on the more complex conditions in interactions while \\citet{jon10} examined how well isolated galaxies were reproduced. We concur with their assessment that this simulation process yields realistic results, although there remain areas of possible improvement such as a more realistic treatments of processes like star formation at smaller scales. Regarding the issue of cold dust, their Figure 16, for example, showed clearly that the observed MIPS 70\\um~and 160\\um~data points lie well away from the simulations colors, indicating the incomplete treatment of cold dust in the simulation. Our PACS and SPIRE data similarly demonstrates (e.g. Figure \\ref{int_fits}) that very often the FIR-submillimeter points are poorly reproduced, although our fits are much closer than the ones in the \\citet{jon10} paper. We compared the dust using two treatments of sub-resolution ISM structure in the simulations; we found that some galaxies match better with one treatment, while others are better matched with the other. With the current simulations, the sub-resolution structure of the ISM remains a significant modeling uncertainty (see Figures \\ref{sunrise} and \\ref{stdev_sunrise} and the associated discussion). However, in future work, this uncertainty can be reduced considerably through the use of higher-resolution merger simulations that include a more-realistic multiphase ISM \\citep[e.g.,][]{hop13a, hop13b}.  This study has demonstrated that with even a relatively modest number of simulations, it is possible to match the SEDs of observed interacting galaxies with one or more mock SEDs from the simulations. Furthermore, the best matches tend to come from relatively constrained regions of the physical parameter space, which suggests that the fits are reasonably non-degenerate. While a more expansive library of simulations may yield better matches to some of our observed systems, these comparisons, together with \\citet{jon10},  demonstrate the feasibility of directly inferring physical quantities of galaxies by direct comparison with forward-modeled mock SEDs rather than through traditional SED modeling, assuming that the simulations span a sufficiently large parameter space with sufficiently fine sampling. Although creating the template SEDs using this forward-modeling approach is orders-of-magnitude more computationally expensive than standard SED modeling, the advantage is that this approach is predictive in the sense that the SEDs are calculated self-consistently from hydrodynamical simulations; the predictive nature of the forward-modeling approach yields additional value beyond inferring physical parameters of the galaxies, such as the possibility of examining the potential evolution and past of these systems.  This work has utilized only the integrated SEDs of the simulated galaxies, but these represent only a tiny fraction of the information available from the simulations. In future work, we will more fully utilize the wealth of data available from the simulations by comparing the simulated and observed galaxy morphologies, for example, and determining what morphological information is encoded in the integrated SEDs."
    },
    "1402/1402.5960_arXiv.txt": {
        "abstract": "Numerical studies of primordial star formation suggest that the first stars in the universe may have been very massive.  Stellar models indicate that non-rotating Population III stars with initial masses of  $140-260\\,\\Msun$ die as highly energetic pair-instability supernovae.  We present new two-dimensional simulations of primordial pair-instability supernovae done with the \\CASTRO{} code.  Our simulations begin at earlier times than previous multidimensional models, at the onset of core contraction, to capture any dynamical instabilities that may be seeded by core contraction and explosive burning.  Such instabilities could enhance explosive yields by mixing hot ash with fuel, thereby accelerating nuclear burning, and affect the spectra of the supernova by dredging up heavy elements from greater depths in the star at early times.  Our grid of models includes both blue supergiants and red supergiants over the range in progenitor mass expected for these events.  We find that fluid instabilities driven by oxygen and helium burning arise at the upper and lower boundaries of the oxygen shell $\\sim$ 20 - 100 seconds after core bounce. Instabilities driven by burning freeze out after the SN shock exits the helium core.  As the shock later propagates through the hydrogen envelope, a strong reverse shock forms that drives the growth of Rayleigh--Taylor instabilities.  In red supergiant progenitors, the amplitudes of these instabilities are sufficient to mix the supernova ejecta. ",
        "introduction": "Modern cosmological simulations suggest that the first stars formed in small pre-galactic structures known as cosmological halos with masses of $\\sim \\Ep6\\,\\Msun$ at $z \\sim$ 20 \\citep{bcl99,abn00, abn02,bcl02,nu01,on07,y08,karlsson2013}.  The original models suggested that Population III (Pop III) stars were 100 - 500 \\Ms\\ and formed in isolation, one star per halo.  Simulations have since shown that some Pop III stars formed in binaries \\citep{turk2009} or small multiples \\citep{stacy2010,get11,stacy2012, get12}.  These and other calculations \\citep{hos11,hir13} indicate that Pop III stars were likely 40 - 500 \\Ms. The properties of primordial stars are key to understanding early cosmological reionization \\citep{whalen2004,kitayama2004,abel2007} and chemical enrichment \\citep{mbh03,ss06,schn06, ss07,bsmith09,ritt12,ss13,cooke2014}.  These stars also populated the first galaxies \\citep{jgb08,get08,jeon11,pmb11, wise12,pmb12, jeon2013} and may be the origin of supermassive black holes \\citep[SMBHs;][]{milos09,awa09, pm11,jlj12a,wf12,agarw12,jet13,latif13c,latif13a,jet13a,wet13b,wet13a}.  The final fate of the first stars depends on their masses.  In particular, non-rotating 140 - 260 \\Ms\\ Pop III stars are thought to die in highly energetic pair-instability supernovae \\citep[PSNe;][]{bark1967, glatzel1985,heger2002,heger2010} \\citep[rotation can extend this lower limit down to 85 \\Ms;][]{cw12}. Stars in this mass range reach central temperatures above 10$^9$ K at densities below $\\Ep6\\,\\gcc$, which favors the creation of electron-positron pairs at the expense of thermal pressure support in the core.  The adiabatic index, $\\gamad$, in the core falls below the critical value of $\\frac43$, {\\color{black} causing it to contract}. Temperatures and densities in the core rise steeply, triggering explosive burning of oxygen and silicon.  The energy release (up to 10$^{53}$ erg) reverses the core contraction and, in most cases, completely unbinds the star, leaving no compact remnant behind \\citep[but see][]{wet13e}. Such events also produce up to 50 \\Ms\\ of \\Ni{}.  Recent events have rekindled interest in PSNe.  PSN are promising candidates for luminous SNe recently  discovered in the local universe, SN 2007bi at $z =$ 0.123 \\citep{2007bi} and SN 2213 - 1745 at $z =$ 2.05 \\citep{cooke12}.  {\\color{black} However, these events can be also explained by other models such as magnetar spin-down \\citep{kasen2010, woosley2010, dessart2012} or H-poor SN ejecta circumstellar interaction \\citep{chat2012a, moriya2013}. The nature of these transits is  still under debate.}   Stars with masses above the canonical limit of 150 \\Ms\\ have also been found, including some with masses greater than 300 \\Ms\\ \\citep{humphrey1979,davidson1997,r136}.  New studies have also shown that PSNe will be visible in the near infrared (NIR) at $z \\sim$ 15 - 20 to the {\\it James Webb Space Telescope} ({\\it JWST}), the Wide-Field Infrared Survey Telescope (WFIRST) and the next generation of extremely large telescopes \\citep{kasen2011,wet12a,hum12,pan12a, wet13c,wet12b,wet13d,ds13,ds14}.  PSNe could therefore probe the masses of the first generation of stars \\citep[see also][]{mw12,wet12c,wet12e,wet12d,chat2012a,mes13a,chen14a}.  Most PSN models to date have been one-dimensional (1D) \\citep{ober1983,stringfellow1988,scann05, dessart2013}.  In the initial stages of a supernova, however, spherical symmetry can be broken by fluid instabilities that require multidimensional simulations.  Recent two-dimensional (2D) models of Pop III PSNe by \\citet{candace2011} found that in most cases either no instabilities or only mild ones arose in the explosion.  But their models proceeded from \\KEPLER{} profiles in which explosive burning was already complete.  In reality, instabilities could be seeded at earlier times by core contraction and nuclear burning.  They could alter the energetics and chemical yields of the explosion itself by mixing hot ash with fuel and enhancing burning.  \\citet{chen2011} examined initial core contraction and bounce in PSNe in 2D but did not evolve the shock to breakout.  {\\color{black} \\citet{chatz2013} modeled the PSNe of rotating stars with different masses in 2D.  However, due to the rotation and metallicity effects, their study was limited to PSNe of blue supergiants.}  Here, we consider the PSNe of non-rotating blue supergiants and red supergiants from 150 - 250 \\Ms\\ in 2D to investigate the formation of dynamical instabilities during the explosion and their impact on energetics and elemental yields.  We describe our numerical methods and progenitor models in Section 2.  The explosions, dynamical instabilities and internal mixing are examined in Section 3.  We conclude in Section 4. ",
        "conclusions": "Our simulations of Pop III PSNe are the first to follow core contraction, nuclear burning, explosion and the end of mixing in 2D in both red and blue supergiants.  As mentioned earlier, instabilities seeded during core contraction and explosive burning in principle can alter the yield of the SN by mixing hot ash with fuel and accelerating burning, unlike instabilities that develop in the reverse shock at later times after the end of burning. Although instabilities do appear at these early stages, they do not affect nucleosynthesis or energy generation for two reasons.  First, most of the explosive silicon burning occurs at the very center of the star, not at the base of the oxygen-burning shell where the instabilities form, so there is no evidence of mixing of hot ash in the core.  Second, burning time scales are so short compared to dynamical times that instabilities do not have time to become large enough to dredge material up from the core and drive mixing.  Later, when the SN shock reaches the helium layer, it burns \\He{} into \\Cx, \\Ox, and \\Ni.  The energy released creates pressure gradients opposite to the gradients in density and mass fractions, causing the formation of RT instabilities that also, in principle, could accelerate helium burning.  But these instabilities do not survive for long because temperatures behind the shock rapidly fall as the star expands. About 100 seconds after core bounce, the postshock gas has cooled to $2\\times\\Ep8\\,\\temp$, and nuclear burning becomes too weak to drive further instability growth.  They then become frozen in mass coordinate until the reverse shock forms.  Although they may briefly enhance helium burning, it does not affect the speed of the shock or its energy.  Instabilities are too weak to dredge \\Ni\\ up from the core at early times in PSNe in either red or blue supergiants. Later, when instabilities with much larger amplitudes form in red supergiants because of the reverse shock, they still do not transport much \\Ni\\ up from lower depths and most of the energy from its decay is deposited locally in the ejecta.  Our simulations therefore suggest that internal mixing in these events will probably not be visible in their observational signatures, in either the order that lines from metals appear in the spectra over time or the appearance of gamma rays from \\Ni\\ decay. The conclusions about the detectability of PSNe at high $z$ derived from prior 1D radiation hydrodynamics calculations therefore still hold.  We note that if these explosions occur at high redshift, neither X-rays and hard UV from shock breakout nor the leakage of gamma rays from \\Ni\\ at later times would be detected because of absorption by the neutral intergalactic medium and the outer layers of the Galaxy \\citep{wet12a,wet12b}.  We find that instability growth in Pop III PSNe is generally much weaker than in 15 - 40 \\Ms\\ Type II SNe, which exhibit rampant mixing \\citep{chevalier1976,fryxell1991,herant1994,candace2010}.  PSN shocks form at the edge of the carbon/oxygen core, which contains 40\\% of the mass of the star, and the envelope beyond it is not dense enough to foster as strong a reverse shock as in less massive stars.  Unlike most PSNe, core-collapse SNe also exhibit fallback onto a central compact remnant that can enhance mixing.  Stellar rotation can reduce the progenitor masses of PSNe \\citep{cw12,chatz2013} but also tends to create blue supergiants in which there is less mixing. This point is important because recent cosmological simulations suggest that many Pop III stars may have been born with rotation speeds close to the breakup value \\citep{stacy11b,get12,stacy13}."
    },
    "1402/1402.2292_arXiv.txt": {
        "abstract": "During hard X-ray observations of the Norma spiral arm region by the {\\em Nuclear Spectroscopic Telescope Array (NuSTAR)} in 2013 February, a new transient source, NuSTAR~J163433--4738.7, was detected at a significance level of 8-$\\sigma$ in the 3--10\\,keV bandpass.  The source is consistent with having a constant {\\em NuSTAR} count rate over a period of 40\\,ks and is also detected simultaneously by {\\em Swift} at lower significance.  The source is not significantly detected by {\\em NuSTAR}, {\\em Swift}, or {\\em Chandra} in the days before or weeks after the discovery of the transient, indicating that the strong X-ray activity lasted for between $\\sim$0.5 and 1.5 days.  Near-IR imaging observations were carried out before and after the X-ray activity, but we are not able to identify the counterpart.  The combined {\\em NuSTAR} and {\\em Swift} energy spectrum is consistent with a power-law with a photon index of $\\Gamma = 4.1^{+1.5}_{-1.0}$ (90\\% confidence errors), a blackbody with $kT = 1.2\\pm 0.3$\\,keV, or a bremsstrahlung model with $kT = 3.0^{+2.1}_{-1.2}$\\,keV. The reduced-$\\chi^{2}$ values for the three models are not significantly different, ranging from 1.23 to 1.44 for 8 degrees of freedom.  The spectrum is strongly absorbed with $N_{\\rm H} = (2.8^{+2.3}_{-1.4})\\times 10^{23}$\\,cm$^{-2}$, $(9^{+15}_{-7})\\times 10^{22}$\\,cm$^{-2}$, and $(1.7^{+1.7}_{-0.9})\\times 10^{23}$\\,cm$^{-2}$, for the power-law, blackbody, and bremsstrahlung models, respectively.  Although the high column density could be due to material local to the source, it is consistent with absorption from interstellar material along the line of sight at a distance of 11\\,kpc, which would indicate an X-ray luminosity $>$$10^{34}$\\,erg~s$^{-1}$.  Although we do not reach a definitive determination of the nature of NuSTAR~J163433--4738.7, we suggest that it may be an unusually bright active binary or a magnetar. ",
        "introduction": "Hard X-ray surveys of the Galaxy provide an opportunity to discover populations of extreme sources.  The promise of such surveys has been partially realized with the {\\em International Gamma-Ray Astrophysics Laboratory} \\citep[{\\em INTEGRAL},][]{winkler03}, which carried out a 20--100\\,keV survey of the entire Galactic Plane and has discovered hundreds of new sources \\citep{bird10,krivonos12}, including new types of High Mass X-ray Binaries (HMXBs), pulsar wind nebulae (PWNe), and magnetic Cataclysmic Variables (CVs).  The {\\em Nuclear Spectroscopic Telescope Array} \\citep[{\\em NuSTAR},][]{harrison13}, which launched in 2012 June and covers the 3--79\\,keV bandpass, is the first focusing hard X-ray telescope in orbit.  While it has a much smaller field of view than {\\em INTEGRAL}, it has much lower background and greatly improved sensitivity.  Thus, one of {\\em NuSTAR}'s science goals is to extend our view deeper into the Galactic plane to look for hidden hard X-ray populations.  During its first year of operation, {\\em NuSTAR} has initiated surveys of $\\sim$1 deg$^{2}$ areas both in the Galactic Center \\citep[e.g.,][]{mori13,nynka13} and in a region which samples the spiral arm population in order to probe potentially different environments where X-ray binaries are found.  A spiral arm region centered on Galactic coordinates of $l = 337.5^{\\circ}$ and $b = 0^{\\circ}$ was chosen for having the highest known density of OB star associations \\citep{russeil03} and HMXBs \\citep{bodaghee07,bodaghee12}. This is part of the ``Norma'' spiral arm region, which was identified early in the {\\em INTEGRAL} mission as having an unusually high density of hard X-ray sources \\citep{tomsick04_munich,dean05,lutovinov05a}.  The combination of active star formation and evidence that compact objects have already formed suggests that a survey by {\\em NuSTAR} may uncover compact objects associated with populations of massive stars such as magnetars or faint HMXBs that are early in their evolutionary process and may have neutron star or black hole accretors.  The full surveys of the Galactic Center and the Norma region will be carried out over a period of $\\sim$2 years; here we report on the discovery of a transient source made during the first part of the survey.  In the following, \\S\\,2 describes observations made with {\\em NuSTAR}, {\\em Swift}, and {\\em Chandra}, as well as the procedures we used to reduce the data.  The results are presented in \\S\\,3, and we discuss possibilities for the nature of the transient in \\S\\,4. ",
        "conclusions": "In considering the nature of NuSTAR~J163433--4738.7, it is useful to estimate the X-ray luminosity that the source reached during its outburst.  Although the distance is highly uncertain, the strong absorption either indicates a large distance through a region of the Galaxy with heavy extinction or absorption local to the source.  The former is a strong possibility since the transient is in the direction of a region of the Galactic plane that is crowded with H~II/molecular cloud regions.  NuSTAR~J163433--4738.7 is at $l = 336.787^{\\circ}, b = +0.014^{\\circ}$, near a group of H~II regions (336.732+0.072, 336.840+0.047, 336.9--0.1, 337.147--0.181, and 337.3--0.1) at a distance of $10.9\\pm 0.2$\\,kpc \\citep{georgelin96,russeil03}\\footnote{This kinematic distance is estimated using the mean Galactic rotation curve, and the uncertainty does not account for possible deviations relative to the mean.}. If the absorption is interstellar, then NuSTAR~J163433--4738.7 is very likely beyond or in these molecular clouds, indicating a 2--10\\,keV unabsorbed luminosity limit of $>$$9\\times 10^{33}$\\,erg~s$^{-1}$ for the blackbody spectrum and $>$$1.4\\times 10^{34}$\\,erg~s$^{-1}$ for bremsstrahlung.  In the following, we consider both possibilities: a distance of $\\sim$11\\,kpc and a peak luminosity of $\\sim$$10^{34}$\\,erg~s$^{-1}$; and a smaller distance with absorption local to the source and a lower luminosity.  Considering the first possibility, a luminosity as high as $10^{34}$\\,erg~s$^{-1}$ would be unprecedented for at least two common types of X-ray sources in the Galaxy. Non-magnetic CVs have strong optical outbursts along with persistent and transient X-ray emission.  However, extensive studies have shown that their X-ray luminosities are in the $10^{29-32}$\\,erg~s$^{-1}$ range \\citep{baskill05,kuulkers06}.  While magnetic CVs can reach higher luminosities \\citep{kuulkers06}, they do not typically show outbursts.  Secondly, active binaries, including RS~CVn systems and low-mass flare stars, produce X-ray flares, some of which can last for about a day. Superflares that peak at X-ray luminosities of $10^{32-33}$\\,erg~s$^{-1}$ have been seen \\citep{franciosini01,osten07,osten10,ps12}, and there have been flares that have released $\\gsim$$10^{37}$\\,erg \\citep{franciosini01}.  If NuSTAR~J163433--4738.7 is at $\\sim$11\\,kpc, then in addition to having a higher peak luminosity, the energy released is $\\gsim$$4\\times 10^{38}$\\,erg, which is based on the source being at its peak luminosity for $\\gsim$40\\,ks (see Figure~\\ref{fig:lc}).  In summary, while we do not rule out the active binary possibility (and see below for further discussion on this topic), at the 11\\,kpc distance, the event seen by {\\em NuSTAR} would need to be extreme for this explanation to be correct.  A source type that might be a good match to the NuSTAR~J163433--4738.7 properties is the class of highly magnetic isolated neutron stars: magnetars.  While these sources are best known for their very bright and brief ($\\sim$0.1\\,s) X-ray and gamma-ray flares, they also have persistent but variable emission that can easily reach $10^{34}$\\,erg~s$^{-1}$ or higher \\citep{wt06}.  The X-ray spectrum of their persistent emission is dominated by a $\\sim$0.5--1\\,keV blackbody \\citep{wt06,mori13}. Furthermore, as mentioned in \\S\\,1, magnetars are associated with regions where high-mass star formation is occurring, such as the Norma region, and there is a known magnetar, SGR~1627--41, that is $\\sim$0.2$^{\\circ}$ from NuSTAR~J163433--4738.7. However, one possible counter-argument to the magnetar hypothesis is that magnetar periods of activity usually last for months.  From Figure~\\ref{fig:lc_overview}, the peak activity from NuSTAR~J163433--4738.7 could not have lasted for more than $\\sim$1.5\\,days based on the {\\em Swift} upper limits, and {\\em Chandra} and {\\em NuSTAR} place a tight upper limit on activity $\\sim$3 weeks after the {\\em NuSTAR} detection.  While we do not detect pulsations or $\\sim$0.1\\,s flares, which would prove that the source is a magnetar, the pulsation search is limited by the statistical quality of the data, and flaring episodes are relatively rare.  A blackbody spectrum could also be produced from an optically thick accretion disk in a black hole binary.  Most often, such sources show temperatures of $\\sim$1\\,keV in the inner parts of their accretion disks when they are at high luminosity $\\gsim$$10^{36-37}$\\,erg~s$^{-1}$, which would require a very large distance of $\\gsim$100--300\\,kpc for NuSTAR~J163433--4738.7. Two sources (GRS~1758--258 and 1E~1740.7--2942) have shown fainter soft states, but their blackbody spectra have been at temperatures of 0.4\\,keV \\citep{smith01} and 0.7\\,keV \\citep{delsanto05}, which are lower than the 1.2\\,keV we see for NuSTAR~J163433--4738.7.  Also, transient black hole binaries typically have outbursts that last for several months, which is very different from NuSTAR~J163433--4738.7.  Considering the second possibility that the source is closer than $\\sim$11\\,kpc, then the high column density must be due to material local to the source. In fact, {\\em INTEGRAL} has found a large number of obscured HMXBs that have column densities of $10^{23-24}$\\,cm$^{-2}$ due to the wind from their supergiant companions \\citep[e.g.][]{walter06}.  While many of the sources in this class are persistent, the Supergiant Fast X-ray Transients (SFXTs), which have some members that are also obscured HMXBs, have outbursts that can be as short as a few hours, and it would not be unusual for an SFXT to produce X-ray emission that lasts for half a day \\citep{smith98a,negueruela06,romano11}.  However, a possible problem for this interpretation is that the spectrum of NuSTAR~J163433--4738.7 is softer than usually seen for active SFXTs.  During flares, IGR~J17544--2619 and IGR~J16479--4514 show a power-law spectrum with a photon index near $\\Gamma = 1.4$ \\citep{rampy09,sidoli13}.  Somewhat softer spectra with $\\Gamma\\sim 2.3$--2.4 can be seen during periods of weaker activity, but their spectra are only as soft as NuSTAR~J163433--4738.7 when they are in quiescence.  For CVs and active binaries, the challenge is to explain a column density as high as we measure with absorption local to the systems.  One possibility is that material expelled from the system (either related to the evolutionary state of the stellar component or related to the X-ray flaring event) could cause increased column density.  However, for one active binary (AX Ari), the column density was seen to increase to $N_{\\rm H} = 1.1\\times 10^{20}$\\,cm$^{-2}$ during an event \\citep{franciosini01}, which is 2--3 orders of magnitude less than is required for NuSTAR~J163433--4738.7.  Variation in optical brightness of some stars has also been attributed to ejection of material from the stars. An example of this is a change of 0.5 magnitudes in the optical seen for KU Cyg \\citep{tang11}.  However, this would translate into a column density near $10^{21}$\\,cm$^{-2}$, which is still far less than is required.  The near-IR information that we have provides only weak constraints on the source type.  Given that we were not able to identify the counterpart, we can only say that the source must be fainter (before and after the X-ray flare) than the brightest source in the error circle (2MASS~J16343288--4738393). This corresponds to $K_{s} > 12.3$, $H > 13.1$, and $J > 14.6$, and these values are consistent with all the possibilities for the source type discussed above.  However, in the scenario where NuSTAR~J163433--4738.7 is a nearby source with local absorption, these limits do provide some constraint.  For example, a very nearby and luminous obscured HMXB is ruled out.  In summary, NuSTAR~J163433--4738.7 is a fast X-ray transient, which has a thermal spectrum with relatively high temperature ($kT = 1.2$\\,keV for a blackbody or 3.0\\,keV for bremsstrahlung).  If its high column density is due to interstellar material, the source is probably distant ($\\gsim$11\\,kpc), making the peak luminosity $\\gsim$$10^{34}$\\,erg\\,s$^{-1}$.  We discuss the origin of the flare and suggest that the most likely possibilities if the source is distant are an unusually bright flare from an active binary or a short outburst from a magnetar.  We also consider the possibility that the source is closer and that the absorption is local to the source.  More {\\em NuSTAR} observations in the Galactic Plane will determine whether such transients are common and, hopefully, shed light on the nature of NuSTAR~J163433--4738.7."
    },
    "1402/1402.6772_arXiv.txt": {
        "abstract": "A drift mode accelerometer is a precision device that overcomes the much of the acceleration noise and readout dynamic range limitations of traditional electrostatic accelerometers. It has the potential of achieving acceleration noise performance of drag-free systems over a restricted frequency band without the need for external drag-free control or spacecraft propulsion. Like traditional accelerometers, the drift mode accelerometer contains a high-density test mass surrounded by an electrode housing, which can control and sense all six degrees of freedom of the test mass. Unlike traditional accelerometers, the suspension system is operated with a low duty cycle so that the limiting suspension force noise only acts over brief, known time intervals, which can be neglected in the data analysis. The readout is performed using a laser interferometer which is immune to the dynamic range limitations of even the best voltage references typically used to determine the inertial acceleration of electrostatic accelerometers. The drift mode accelerometer is related to the like-named operational mode of the LISA Pathfinder spacecraft, which will be used to estimate the acceleration noise associated with the LISA Pathfinder front end electronics. This paper describes operation of such a device, develops models for its performance with respect to satellite geodesy and gravitational wave astrophysics applications, and discusses methods for testing its performance using torsion pendula in the laboratory and the LISA Pathfinder mission in space. ",
        "introduction": "Precise measurement of inertial acceleration is vital to many space-borne gravitational science missions, including satellite geodesy \\cite{GRACE}, fundamental physics experiments \\cite{microscope, prl} and gravitational wave observation \\cite{LISA}. The most precise accelerometers manufactured to date are the electrostatic accelerometers produced by ONERA, which are capable of measuring spacecraft acceleration relative the the inertial frame  to $\\sim 10^{-11} \\ \\mathrm{m/sec^2 Hz^{1/2}}$ from roughly 1 mHz to 1 Hz \\cite{touboul2012}. These accelerometers have been used for Low-low Satellite-to-satellite tracking missions including GRACE \\cite{GRACE} and for gravity gradiometer missions such as GOCE \\cite{GOCE}.  These instruments are comprised of an internal free-floating metallic test mass that is surrounded by an electrode housing. The electrodes on the internal surface of the housing both sense the test mass' position capacitively and actuate it via electrostatic forces. The position measurement is used to drive the electrostatic suspension system to keep the test mass centered in its housing. The inertial acceleration of the spacecraft is proportional to the suspension force applied to the test mass to keep it centered.  Electrostatic accelerometers are limited by two inter-related factors: 1) suspension force noise and 2) acceleration measurement noise. Both are ultimately related to the stability of voltage references, where the current state of the art is $\\sim 2 \\times 10^{-6}$ \\cite{GOCEaccel}. For the application of Earth geodesy the low frequency acceleration of a low Earth orbiting satellite can be as high as $\\sim 10^{-5} \\ \\mathrm{m/sec^2}$. Therefore the resulting acceleration noise on the test mass due to the suspension system is at least $2 \\times 10^{-11} \\ \\mathrm{m/sec^2}$. Since the applied suspension force \\emph{is} the acceleration measurement, the acceleration measurement noise would be on this same order. To improve accelerometers significantly beyond the $10^{-11} \\ \\mathrm{m/sec^2}$ level, the suspension force noise must be removed \\emph{and} the sensor used to measure acceleration must be changed.  Drag-free technology, conceived of in the 1960's \\cite{Lange, debra}, has been the most promising approach to breaking through these acceleration noise limits. Two drag-free approaches have been demonstrated on three separate missions. The first is an ``accelerometer-mode'' drag-free, where an electrostatic accelerometer is used as the primary sensor and a propulsion system is used to counter the drag-force acting on the satellite so that the nominal test mass suspension force is reduced. The spacecraft acceleration measurement is still limited by voltage reference stability, but the nominal voltage applied to the housing electrodes is reduced, therefore the electrostatic force noise is also reduced. Both Gravity Probe B \\cite{prl} and the GOCE \\cite{GOCE} missions operated in accelerometer-mode drag-free. Using this approach Gravity Probe B achieved an acceleration noise of $4 \\times 10^{-11} \\ \\mathrm{m/sec^2 Hz^{1/2}}$ \\cite{bencze} and GOCE achieved a differential acceleration noise measurement between test masses accurate to $\\sim 10^{-12} \\ \\mathrm{m/sec^2 Hz^{1/2}}$ in the 1 mHz to 1 Hz frequency band \\cite{GOCEaccel}.  The other drag-free operating mode is 'true' drag-free, where the suspension force is turned completely off, at least in one degree of freedom. Triad I with its DISturbance COmpensation System (DISCOS) operated in this manor \\cite{discos1, discos2}, as will the Laser Interferometer Space Antenna (LISA) in the future. A fundamental difference between accelerometers and true drag-free is that the basic measurement for a true drag-free system is displacement variations, instead of acceleration variations. Of course one can always convert displacement to acceleration and vice-versa.  A drift-mode accelerometer (DMA) as defined here is a traditional electrostatic accelerometer where the test mass suspension force is operated with a low duty cycle. Larger suspension forces are used, but over a much shorter period of time so that the average suspension force is the same as that of a traditional accelerometer. By switching the suspension system on and off with a constant frequency and low duty cycle ($< 0.1$), the suspension system force noise is restricted to known, short intervals, which repeat with a frequency chosen to be above the science frequencies of interest.  Cycling the suspension system eliminates suspension force noise while the suspension system is off, but there still is the problem of precisely measuring the inertial acceleration of the satellite in the presence of a large zero-frequency acceleration. Here, laser interferometry provides the solution. In most upcoming precision gravity missions, the measurement of interest is the relative displacement (or acceleration) between two or more inertially fixed test masses. GRACE Follow-on, GRACE-II, and LISA \\cite{LISA} are all examples. In all of these missions, the laser interferometer system already exists and is used to measure range variations between spacecraft. If an interferometer is used to also measure distance variations between a reference point on the spacecraft and the test mass, then this measurement can be used to estimate the inertial acceleration of the spacecraft, assuming that the test mass can be treated as inertially fixed over the short interval when the suspension system is off. Second order finite differencing provides the simplest method. Although other approaches discussed in this paper can provide more accurate estimates.  Laser interferometers have been demonstrated with extremely large dynamic range. The LISA Interferometric Measurement System for example can measure pm variations over 1000 sec between spacecraft that have relative velocities of 10 m/sec. This represents a dynamic range of $10^{22}$.  The name drift mode is taken from an operating mode of LISA Pathfinder (LPF) \\cite{LPF, LPFdriftmode}. LPF contains two free-floating test masses. The spacecraft can only fly drag-free about one of them (naturally) and, therefore, the other test mass must be suspended against the gravity gradient (and other) forces which act upon it. In order to assess the acceleration noise associated with the suspension electronics, the drift mode was conceived. The suspension system is turned on and off with a low duty cycle (1 sec on and 200 sec off). In between ``kicks'' the test mass follows approximate parabolic trajectories, when measured relative the the other test mass. These parabolas are fit to second order polynomials and the fit residuals are used to calculate variations in the differential acceleration between the two test masses. Since the goal of the drift mode for LPF is simply to determine the acceleration noise on the test masses due to the actuation electronics, the time between kicks was chosen to be relatively long (200 sec). The interferometer data during the kicks is discarded and replaced with a model of the acceleration noise that makes various assumptions about the nature of the noise \\cite{LPFdriftmode}. In contrast, for the DMA we wish to make no assumptions about the inertial acceleration of the satellite and therefore, we choose a kicking frequency that lies above the science signal of interest. ",
        "conclusions": "The drift mode accelerometer is a modified electrostatic accelerometer potentially capable of acceleration noise performance similar to that of drag-free systems without the need for drag-free control or associated precision propulsion. A DMA consists of a dense test mass that is freely floating inside an electrode housing, which can both sense its position capacitively and actuate it electrostatically. Unlike traditional electrostatic accelerometers, the suspension system is operated with a low duty cycle and with a cycling frequency that is chosen to be above the science signals of interest. Measurement of spacecraft acceleration is made using a laser interferometer, which is not limited by dynamic range. Two applications of the DMA, Earth geodesy and gravitational wave observation, are studied. Both represent gravitational science missions where the DMA might be used to replace drag-free operation. For gravitational wave observation, the combination of the existing LISA Pathfinder gravitational reference sensor and the LISA local (short-arm) interferometer can be operated as a drift mode accelerometer, with acceleration noise performance close to that required for LISA. Detailed modeling and analysis is still required to fully determine the acceleration noise performance and instrument requirements and constraints. Laboratory testing using torsion pendula provide one promising approach for demonstrating the performance and operation of the drift mode accelerometer."
    },
    "1402/1402.1687_arXiv.txt": {
        "abstract": "We present CMDs in the $V$ and $I$ bands for fifteen star clusters in the Small Magellanic Cloud (SMC) based on data taken with the Very Large Telescope (VLT, Chile). We selected these clusters from our previous work \\citep{par09}, wherein we derived cluster radial velocities and metallicities from Calcium II infrared triplet (CaT) spectra also taken with the VLT. We discovered that the ages of six of our clusters have been appreciably underestimated by previous studies, which used comparatively small telescopes, graphically illustrating the need for large apertures to obtain reliable ages of old and intermediate-age SMC star clusters. In particular, three of these clusters, L\\,4, L\\,6 and L\\,110, turn out to be amongst the oldest SMC clusters known, with ages of 7.9 $\\pm$ 1.1, 8.7 $\\pm$ 1.2 and 7.6 $\\pm$ 1.0 Gyr, respectively, helping to fill a possible ``SMC cluster age gap'' \\citep{gla08a}. Using the present ages and metallicities from \\citep{par09}, we analyze the age distribution, age gradient and age metallicity relation (AMR) of a sample of SMC clusters measured homogeneously. There is a suggestion of bimodality in the age distribution but it does not show a constant slope for the first 4 Gyr \\citep{pia11b}, and we find no evidence for an age gradient. Due to the improved ages of our cluster sample, we find that our AMR is now better represented in the intermediate/old period than that we derived in \\citealt{par09}, where we simply took ages available in the literature. Additionally, clusters younger than $\\sim$ 4 Gyr now show better agreement  with the bursting model of \\citet{pag98}, but we confirm that this model is not a good representation of the AMR during the intermediate-age/old period. A more complicated model is needed to explain the SMC chemical evolution in that period. ",
        "introduction": "As it is widely known, the Small Magellanic Cloud (SMC) is the closest dwarf irregular galaxy to the Milky Way. For this reason, it is the ideal local analog to this type of primeval distant galaxies. Due to its low metallicity, the SMC constitutes an excellent field of study to explore the role played by metallicity in these galaxies. Star clusters of the Magellanic Clouds (MCs), particularly those of the SMC, are important to trace their chemical evolution and star formation history. Because of the richness and variety of the SMC star clusters, they are also important testbeds for theoretical models of stellar evolution at intermediate age and moderately low metallicity (e.g., \\citealt{fer95}). Besides, the SMC clusters have been used as empirical templates to interpret the unresolved spectra and colors of distant galaxies (e.g., \\citealt{betal02}).  In order to study the early star formation history of a galaxy, it is crucial to characterize well the oldest stellar populations. Globular Clusters (GCs) are ideal objects as tracers of  the oldest populations. In addition, although the GC system of a galaxy exhibits an age range, the oldest GCs clearly belong to the oldest stellar systems that were formed in the early Universe and still survive. In those galaxies where the age of the oldest field populations can be determined and compared with the age of the Galactic GCs, it was found that, in general, the age difference between these  two populations is not significant within the errors and that the oldest GCs in different galaxies in the Local Group are coeval. This indicates a common  epoch of initial star formation in the Milky Way (MW) and its companions in the Local Group. However, the SMC is the only Local Group galaxy with a significant cluster population which  seems not to share this  common epoch of early cluster formation.  In fact, although this galaxy has a large number of young and intermediate-age clusters, it contains only one cluster, NGC\\,121, confirmed to be older than 10 Gyr and considered the only known SMC GC, in the sense that its age is comparable to those of the bulk of the Galactic GCs.  This has  given rise to a series of speculations and controversies regarding the chemical evolution of the SMC and that of its companion, the Large Magellanic Cloud (LMC).  NGC\\,121, however, is clearly not as old as the oldest GCs of the LMC or of the MW \\citep{gla08a}. In fact, NGC\\,121 has an estimated age of only 10.5 Gyr \\citep{gla08a} instead of 12-12.5 Gyr found for the oldest clusters of the LMC and  MW (e.g \\citealt{dot10}).  Even though recent efforts have uncovered at least one other SMC cluster approaching the age of NGC 121 (HW42, with an age of 9.3 Gyr - \\citealt{pia11a}), it is now clear that the SMC lacks the very oldest GCs that are found in other massive Local Group galaxies. The reason for this is presently unknown.\\\\  The next oldest cluster in the SMC currently known after NGC\\,121 and HW\\,42, L\\,1,  has an age of only 7.5 Gyr, based on deep HST data \\citep{gla08a}. Interestingly, the LMC has a well-known large cluster age gap between 3 and 12 Gyr within which only one cluster exists. We also know, however, that the LMC contains 15 or 16 clusters older than 10 Gyr (see, e.g., \\citealt{dut99}). The SMC indeed also has a potential cluster age gap (as suggested by \\citealt{gla08a}), with only two clusters presently known between 7.5 and 13.7 Gyr, a gap covering almost half the age of the galaxy since the Big Bang.\\\\  At present,  there is no  clear explanation for the lack of clusters in the so-called age gap in the LMC. Likewise, no scenario has been proposed for why there are no truly ancient clusters in the SMC or for its possible cluster age gap, although this latter has not been well studied and the sample size is quite small. The fact that NGC\\,121 is the only known relatively old SMC cluster means that  all conclusions related to the early evolution of this galaxy's cluster system are based only on this single object. It is evident, in this context, that the discovery of other old cluster in the SMC would have important implications related to its early formation and chemical evolution, even clusters inhabiting the possible age gap between NGC 121 and L1. \\\\    Our group began investigating MC clusters using the powerful technique of CaII triplet spectroscopy (CaT) with FORS2 on the VLT several years ago. As shown by \\citet{col04}, CaT is a very efficient and accurate metallicity indicator, with minimal age effects. Our first study \\citep{gro06} yielded excellent data for 28 LMC clusters. We found a very tight metallicity distribution for intermediate age clusters, no metallicity gradient and confirmed that the clusters rotate with the disk. We followed this up with an initial study of SMC clusters \\citep {par09} (hereafter Paper 1). We used FORS2 + VLT to obtain CaT spectra of more than 100 stars in 15 populous SMC clusters, spanning a wide range of ages and metallicities. We derived velocities to a few km s$^{-1}$  per star and mean cluster metallicity (from $>$7 high probability members) to 0.05 dex. We examined the metallicity distribution, metallicity gradient, the age-metallicity relation (AMR) and kinematics for the cluster sample. We found a suggestion of bimodality in the metallicity distribution, no evidence for a gradient and a hint of rotation in our cluster sample, but the kinematics are dominated by the velocity dispersion.\\\\  While the cluster metallicities constituted a reliable as well as homogeneous source of information, since they were determined by applying the same well-proven technique, based on excellent data, cluster ages unfortunately  are on a heterogeneous and relatively poorly-determined scale. This results from the fact that we took the ages available from the literature, which were derived by applying several different methods. These clusters were selected mainly from CMDs in the Washington system published by \\citet{pia05a,pia07b,pia07c} which were based on data obtained on the Cerro Tololo Inter-American Observatory 0.9m telescope. Unfortunately, such data are not optimum for deriving accurate ages for old and intermediate-age SMC clusters, given the small aperture and large pixel scale. The lack of accuracy and homogeneity in the cluster ages limits our confidence in the study of the chemical evolution of the SMC, despite the excellent metallicities in hand. \\\\  In order to improve that situation, we decided to redetermine the ages of the 15 SMC clusters studied in Paper I on a homogeneous scale. As explained there, before obtaining the CaT spectroscopic observations, the VLT staff obtained images of the selected clusters in the $V$ and $I$ bands. These images, usually called {\\it pre-images}, were necessary to build the instrumental color-magnitude diagrams (CMDs) which allowed us to select the cluster red giants that were observed spectroscopically and to perform the requisite astrometry for these targets in order to ensure good slit positioning. {\\it Pre-images} are not generally used for purposes  other than this and thus are usually taken with  the minimum exposure times necessary, in this case to obtain reasonable photometry for the upper red giant branch of each cluster. This unfortunately prohibits using them for something as demanding as determining accurate ages, which requires good data to well below the MS. However, we realized that for a minimum additional cost in overhead we could lengthen the preimages from the usual few second exposures to several minutes. Given the 8-m aperture and generally good to excellent seeing available, this ensured that these data would go much deeper than any available and would yield excellent material from which to derive ages. Thus, we put these preimages to very good astrophysical use! However, in Paper 1 we only used the literature ages as time did not permit us to obtain the required photometry. Also, Paranal Observatory does not provide standard star images in order to determine the transformation of magnitudes and colors to the standard system. Thus, standardizing the preimages was problematic and we waited until we could find an appropriate solution. We have subsequently solved this either by obtaining independent calibration photometry and/or obtaining the ages differentially, without the need for standardization. Hence, the time was ripe for investigating the cluster ages. \\\\  In this Paper, we report on the ages we derive for our cluster sample from the preimages and related results. In Section 2, we detail the procedure used to obtain the uncalibrated photometry. Section 3 describes our initial age determination procedure. Next we present the CMDs of each cluster and investigate two existing age calibrations. Here we find that several of our clusters are substantially older than indicated by previous studies and indeed turn out to be amongst the oldest star clusters known in the SMC. Section 5 describes the importance of such clusters for examining the early chemical and dynamical evolution of the SMC. In addition, we derive a new age calibration appropriate for old to intermediate-age SMC clusters, which yields ages for our oldest clusters supporting their extreme age. In order to further strengthen these results, we also present isochrone ages based on calibrated photometry we derive for the three most interesting old cluster candidates. In Section 6 we discuss our results, while Section 7 summarizes our findings. ",
        "conclusions": "We measured the parameter $\\delta V$ in the $V$ vs. $V-I$ CMD of 15 SMC star clusters and we used the calibrations of \\citet[CC94]{cc94} to derive  their ages. We found that three of them, L\\,110, L\\,4 and L\\,6, are newly discovered very old SMC clusters ( $t =$  7.6 $\\pm$ 1.0, 7.9 $\\pm$ 1.1 and 8.7 $\\pm$ 1.2, respecively).  Consequently,  we have enlarged considerably the number of known old clusters in this galaxy. We corroborate the old nature of these new old clusters, from two independent methods. First, we derived a new calibration (SMC $\\delta$ calibration), using the SMC clusters studied by \\citet{gla08a,gla08b} with deep HST photometry. Secondly, we calibrated our VLT photometry with independent observations and we fitted different sets of isochrones to the calibrated CMD of these old clusters.  Finally, using ages derived from the combination of $\\delta V$ and CC94's calibration, and metallicities from \\citet{par09} we analyzed the age distribution (AD), age gradient and the age metallicity relation of the SMC. Our main results are:  \\begin{itemize}  \\item We combined our 15 SMC cluster with 35 clusters studied by other authors. The AD of this combined sample of 50 clusters reproduces the two peaks found in \\citet{pia11b} at $\\sim$ 2 and $\\sim$ 5 Gyr (corresponding to possible episodes of cluster formation) but it does not show a constant slope in the AD  for the first  $\\sim$ 4 Gyr. The peak at 5 Gyr is in reasonable agreement with the peak in the field star formation rate at about 4.5 Gyr found recently by the HST study of \\citet{wei13} but their second peak at 9 Gyr is not found in our data.  \\item  The extended sample supports the idea of the nonexistence of an age gradient as has been previously suggested by other authors (see, e.g., \\citealt{pia11c,gla08b}). In particular, we find clusters younger than 3.5 Gyrs at a large range of galactocentric distances and not just near the center, as suggested by the field star data of \\citet{wei13}  \\item Combining our cluster sample with those observed by DH98 and Glatt et al. (also with CaT) and the one cluster with high resolution metallicity, we compiled a sample of 25 clusters on a homogeneous metallicity scale with relatively small errors and study the cluster AMR. We corroborate our Paper I result that for ages smaller than $\\sim$ 4 Gyr, the chemical enrichment of the SMC is well represented by the PT98 model but for explaining the intermediate/old SMC period a more complicated model is needed. We now have seven clusters older than 6 Gyr  homogeneously studied so we have doubled the representation of that period.  \\item According to the calibration of CC94, our three new old clusters lie between L\\,1 and NGC\\,121, which suggest that the suggested SMC cluster age gap claimed by \\citet{gla08a} is probably not real.  \\end{itemize}"
    },
    "1402/1402.4777_arXiv.txt": {
        "abstract": "The formation of supermassive Population III stars with masses $\\gtrsim$ 10,000 \\Ms\\ in primeval galaxies in strong UV backgrounds at $z \\sim$ 15 may be the most viable pathway to the formation of supermassive black holes by $z \\sim$ 7.  Most of these stars are expected to live for short times and then directly collapse to black holes, with little or no mass loss over their lives.  But we have now discovered that non-rotating primordial stars with masses close to 55,000 \\Ms\\ can instead die as highly energetic thermonuclear supernovae powered by explosive helium burning, releasing up to 10$ ^{55}$ erg, or about 10,000 times the energy of a Type Ia supernova. The explosion is triggered by the general relativistic contribution of thermal photons to gravity in the core of the star, which causes the core to contract and explosively burn.  The energy release completely unbinds the star, leaving no compact remnant, and about half of the mass of the star is ejected into the early cosmos in the form of heavy elements.  The explosion would be visible in the near infrared at $z \\lesssim$ 20 to {\\it Euclid} and the Wide-Field Infrared Survey Telescope (WFIRST), perhaps signaling the birth of supermassive black hole seeds and the first quasars. ",
        "introduction": "\\lSect{introduction}  The existence of supermassive black holes (SMBHs) in most massive galaxies today poses significant challenges to the paradigm of hierarchical structure formation \\citep{kormendy1995,ferra2000,ferra2005,gebh2000,matte2005,beif2012,mcco2013}.  In particular, it is not known how some SMBHs reached masses of 10$^9$ \\Ms\\ by $z \\sim$ 7, less than a Gyr after the big bang \\citep{fan2002,fan2006,mort11}.  The two leading contenders for the origins of SMBHs are 100 - 500 \\Ms\\ Population III (Pop III) stars at $z \\sim$ 20 and supermassive (10$^4$ - 10$^5$ \\Ms) Pop III stars at $z \\sim$ 10 - 15 \\citep[e.g.,][]{rees1984,madau2001,vol2010, vol12,karlsson2013}.  The main argument against conventional Pop III stars as SMBH seeds is that their BHs must accrete at the Eddington limit continuously down to $z \\sim$ 7 to become supermassive.  How they could sustain such growth is unclear because they form in low-density relic \\HII\\ regions that delay initial accretion \\citep{whalen2004,johnson2007}. Later, when they do accrete, the relatively shallow dark matter potential wells in which they reside cannot retain their fuel supply because of radiative feedback \\citep{awa09}, so they grow at well below the Eddington limit \\citep{pm11,pm12,pm13}.  Furthermore, many low-mass Pop III black holes are ejected from their halos (and thus their fuel supply) at birth by asymmetries in their explosion engines \\citep{wf12}.  For these reasons, there has been growing interest in supermassive stars (SMSs) as candidates for SMBH seeds \\citep[e.g.,][]{loeb1994,bromm2003,bege2006,bromm2011, jet13}.  Such stars could form in atomically-cooled halos at $z \\sim$ 15. In this scenario, a primeval galaxy forms in a strong Lyman-Werner (LW) UV background that sterilizes its constituent halos of H$_2$ prior to assembly \\citep{agarw12,jdk12}, preventing them from forming stars. When the protogalaxy reaches a mass of $\\sim$ 10$^8$ \\Ms\\ and virial temperature of $\\sim$ 10$^4$ K, atomic cooling triggers the rapid collapse of gas at its center at rates of up to 1 \\Ms\\ yr$^{-1}$ \\citep{ln06,spaans06,wta08,rh09,sbh10, latif13c,latif13a,schl13}.  In this manner, a 10$^4$ - 10$^5$ \\Ms\\ Pop III star could be formed in less than a Myr.  The evolution of supermassive protostars and stars has been studied with both analytical and numerical techniques for over 40 years \\citep{fowler1966,wheeler1977, bond1984,carr1984,fuller1986,bk98,baum99,begel10,hos13,schl13}. Most SMSs are thought to become fully convective and then later collapse to massive BHs \\citep[e.g.,][]{st79, shib02,reis13}.  However, previous calculations of SMS evolution have been done at low resolution, considered a relatively coarse grid of models in mass, and in some cases ignored first-order, post-Newtonian general relativistic corrections to gravity. The latter could play an important role in the evolution of 10$^3$ - 10$^5$ \\Ms\\ stars, especially during their pre-explosion phase.  We have revisited the evolution of Pop III SMSs with high-resolution one-dimensional (1D) simulations with updated nuclear reaction rates \\citep{heger2001, heger2002} and post-Newtonian corrections to gravity.  We have found a 55,500 \\Ms\\ star that explodes as a highly energetic thermonuclear supernova (SN) instead of collapsing to a BH \\citep[see also][]{montero12}. Our discovery suggests that there may be a range of high-mass stars that can die as SNe.  In Section 2 we describe our stellar evolution model.  The explosion of the star is examined in 1D in Section 3 and in two dimensions (2D) in Section 4.  We discuss the potential impact of this explosion on its host galaxy and its visibility at high redshift in Section 5. ",
        "conclusions": "\\lSect{gsn_conclusions}  We have discovered that Pop III stars with masses of 55,500 \\Ms\\ may explode as SNe instead of collapsing to BHs. GR effects, rather than the pair instability, trigger the explosion, which at $\\sim$ 9 $\\times$ 10$^{54}$ erg is the most energetic thermonuclear SN known. Mixing in 2D enhances burning and thermonuclear yields during the explosion, which completely unbinds the star and leaves no compact remnant.  Energy release is primarily due to explosive helium burning after the onset of central helium burning. The explosion yields mostly elements between carbon and silicon, with almost no iron group elements.  Besides enhancing burning, mixing during the SN can dredge up heavy elements from deeper layers and cause them to appear at earlier times in spectra. {\\color{black}  \\citet{montero2012} recently studied the relativistic collapse of SMS and discovered that explosions of SMS occurred even for stellar masses up to $5\\times10^{5}$ \\Ms, which are powered by the hydrogen burning due to hot CNO cycle. In their simulations, they use simplified stellar models, with a single fluid of fixed hydrogen, helium, and CNO metal abundance. More importantly, all of their exploded models are associated non zero metallicity, however, there is no metal inside the Pop~III stars when they are born. Hence the hot CNO cycle may not occur for the Pop~III SMS at early on.  Instead, our SMS model evolves from the main sequence through collapse to explosion with the detailed stellar physics which self-consistently produces CNO from the primordial elements. This provides a more realistic pre-supernova model. The differences in pre-supernova models as well as other microphysics used can significantly alter the explosion mechanics that may explain the discrepancies between their findings and ours. }   Our simulations are approximate for several reasons. First, we do not model the evolution of the protostar from much lower masses.  Instead, the star is initialized at the beginning of the main sequence in \\KEPLER{}.  Next, we did not evolve the star under the heavy ongoing accretion that gave birth to the star. We also did not consider stellar rotation, which could lower the mass at which the star explodes \\citep[e.g.,][]{cw12,cwc13}. There is growing evidence that many Pop III stars are born with high rates of spin \\citep{stacy11b,get12}.  Finally, we did not include UV feedback from the star, which could regulate its accretion rates.  However, in some cases luminosity from the star is thought to terminate accretion \\citep{johnson2012}, so our assumption that the star has a constant mass is plausible.  Efforts are now underway to survey Pop III protostellar evolution under a variety of accretion rates and SMS evolution with much larger and finer grids in mass.  What effect would such an explosion have on the protogalaxy that gave birth to the star?  The answer depends on the ambient density of the SN.  If the SMS grows by accretion through a disk its UV radiation will break free of the disk and ionize the surrounding envelope.  In these circumstances the star would explode in low densities, and \\citet{jet13a} show that the SN would drive all the gas from the protogalaxy, perhaps engulfing nearby protogalaxies with metals.  Much of this material would later fall back to the halo on timescales of 50 - 100 Myr and form stars.  If accretion is spherical, then the star will not ionize its envelope and it will explode in very high densities \\citep{johnson2012,hos13,schl13}.  In this scenario, \\citet{wet13b,wet13a} find that much of the energy of the explosion is promptly radiated away by bremsstrahlung X-rays and inverse Compton scattering of cosmic microwave background (CMB) photons, but that the SN remnant still expands to roughly the virial radius of the halo, $\\sim$ 1 kpc, and then collapses back into the halo.  Collapse thoroughly mixes baryons in the halo with metals and may trigger a starburst that would easily distinguish this protogalaxy from its dimmer and less rapidly evolving neighbors.  Radiation hydrodynamical simulations show that the GSN would be visible at $z \\lesssim$ 20 to future all-sky NIR surveys by {\\it Euclid}, the Wide-Field Infrared Survey Telescope (WFIRST) and the Wide-field Imaging Surveyor for High-Redshift (WISH) \\citep{wet12d}. The wide survey fields and high sensitivities of these missions would enable them to detect such an event anywhere in the universe, even if their numbers are small.  Furthermore, a single GSN can create 100 times the chemical yield of a PSN.  But unlike PSNe, which synthesize more iron group elements, GSNe would mainly enrich primordial gas with elements from \\Cx{} to \\Si.  Traces of GSNe might  therefore be found in early galaxies that are \\Fe{} deficient but enhanced with \\Cx{} and \\Ox{} \\citep{keller2014}. Whether in future NIR campaigns or in the fossil abundance record, the detection of GSNe in the early universe may soon signal the birth of SMBH seeds and the first quasars."
    },
    "1402/1402.7068_arXiv.txt": {
        "abstract": "Stellar coronae have been invoked to explain the apparently extragalactic dispersion measures observed in fast radio bursts. This paper demonstrates that the suggested plasma densities would lead to deviations from the standard dispersion curve that are inconsistent with the data. The problem is then turned around and higher-order dispersion terms are connected to the moments of the density distribution along the line of sight. The deviations quantified in three observed bursts are analysed and a lower limit on the maximum electron density is obtained in one case, although with considerable uncertainty. Selection effects are then discussed and shown to be non-restrictive in relation to plasma density, except at the lowest frequencies and highest temperatures. ",
        "introduction": "\\label{section:intro}  Fast radio bursts (FRBs) are a recently discovered class of bright ($\\sim 1\\,\\mathrm{Jy}$) non-repeating radio transients of $\\sim1\\,\\mathrm{ms}$ duration \\citep{lorimeretal2007, keaneetal2012, thorntonetal2013}. The bursts arrival time $t$ vs. observed frequency~$f$ plots adhere to the cold plasma dispersion law, $t\\propto f^{-2}$, very accurately and high values of their dispersion measure, $\\DM\\sim(300-1000)\\,\\mathrm{cm}^{-3}\\,\\mathrm{pc}$ observed far from the Galactic plane suggest an extragalactic origin for FRBs at cosmological distances, $z\\gta0.1$.  The distances, considerable flux densities and short durations imply energies of the FRB sources of up to $10^{40}\\,\\mathrm{erg}$ and brightness temperatures of order $10^{33}\\,\\mathrm{K}$. This can only be achieved by coherent emission with exteme bunching ratios, $q/e\\sim10^{15}$ and involving highly relativistic outflows with Lorentz factor of order $\\Gamma\\gta10^2$ at the lowest \\citep{katz2013}. A number of candidate astrophysical phenomena have been suggested to power the bursts, including magnetar hyperflares \\citep{popovpostnov2007}, supernova impacts on magnetospheres of their companion neutron stars \\citep{egorovpostnov2009}, white dwarf or neutron star mergers ({Kashiyama}, {Ioka} \\& {M{\\'e}sz{\\'a}ros} 2013; Pshirkov \\& Postnov 2010; \\citealt{totani2013}), supramassive neutron star collapse into a black hole \\citep{falckerezzolla2013,zhang2014} and evaporation of primordial black holes \\citep{keaneetal2012}. Recently, \\cite{kulkarnietal2014} have extensively reviewed various proposals singling out the hyperflare model as the most attractive. However, despite the success of these models in predicting the rate, duration and energetics of the observed FRBs, explanations of the emission mechanism, in particular the required bunching ratios and Lorentz factors have so far, with a notable exception of \\cite{lyubarsky2014}, been phenomenological or altogether {\\it ad hoc}.  This has led some authors to question the extragalactic -- or, indeed, extraterrestrial -- origin of the FRB activity. An archival search of the data from the Parkes telescope prompted by \\cite{lorimeretal2007} resulted in the discovery of perytons, a class of apparently terrestrial signals that share some, though not all, of the properties of the original FRB (`the Lorimer burst`), which has cast doubt on its extragalactic interpretation \\citep{burkespolaoretal2011}; recently, perytons have been confirmed as a worldwide phenomenon, not limited to Parkes (Saint-Hilaire, Benz \\& Monstein 2014). Nevertheless, the differences to the Lorimer burst remained while subsequent discoveries of five additional FRBs \\citep{keaneetal2012, thorntonetal2013} which diverge even more from the perytons in their properties seem to have established FRBs as a separate class of transients likely of celestial origin.  Loeb, Shvartzvald \\& Maoz (2014) suggested an alternative interpretation of the high value of the dispersion measure characteristic of FRBs. The authors propose that, rather than being due to a cosmological path length, the observed column density of the electrons is primarily contributed by the high density of electrons, $n_e\\sim10^{10}\\,\\mathrm{cm}^{-3}$ in the coronal plasma of a main-sequence star, which integrates to a $\\DM\\sim300\\,\\mathrm{cm}^{-3}\\,\\mathrm{pc}$ over a stellar scale path length of $10^{11}\\,\\mathrm{cm}$. If an FRB source is located at the base of the corona of a star within $\\sim1\\,\\mathrm{kpc}$ of the Sun, the apparent brigthness, duration and rate of the FRBs might be consistent with the properties of rare, most powerful coherent radio bursts observed at some flaring stars. These bursts are thought to occur via the cyclotron maser mechanism and do not require extreme physical conditions at their sources.  \\cite{luan2014} has criticised the flaring star interpretation by computing the free-free absorption in the coronal plasma. The observed DM implies the absorption that would conceal any radio signal generated below the corona unless it is unrealistically extended or hot. For the burst to remain visible then, it needs to accumulate most of its DM beyond the corona, largely stripping the model of \\cite{loebetal2014} of its explanatory power.  The present paper considers another effect of the high electron densitiy invoked by \\cite{loebetal2014}. % Section~\\ref{section:dispersion} shows that it is high enough for dispersion law terms beyond the standard $f^{-2}$ to become important, and they are not observed. Conversely, Section~\\ref{section:moments} explains how deviations from the standard curve that are visible in the data could constrain the moments of the electron density distribution along the line of sight; estimates available for three FRBs are discussed. Section~\\ref{section:discussion} concludes the paper by discussing the role of selection effects. ",
        "conclusions": "\\label{section:discussion}  In Section~\\ref{section:dispersion} the electron densities $n_e\\sim10^{10}\\,\\mathrm{cm}^{-3}$ suggested by \\cite{loebetal2014} to explain the dispersion measures of FRBs were shown to be inconsistent with the upper limit on the deviation of the observed dispersion curves from the standard, low-density cold plasma law. Densities that are allowed by the data, $n_e\\sim10^{8-9}\\,\\mathrm{cm}^{-3}$, are not extreme on their own -- such is the density at the base of the corona of the Sun -- but require extreme paths, $D\\sim10^2\\,\\mathrm{R}_\\odot$, to integrate to observed DMs, bringing the validity of the stellar model into question.  Section~\\ref{section:moments} turned the problem around and introduced a~simple, moment-based framework with which the deviations of the dispersion law from its low-density limit can be analysed. It was shown that the leading term of the deviation, proportional to the second moment of $n_e$, disguises itself by biasing the estimate of the first moment, the DM, high; this leaves only 4 to 5 per cent of the signal in a form that cannot be so masked. The current surveys are unlikely to probe beyond the second moment and the data available so far are consistent with the standard law to at least $2\\sigma$.  The free-free absorption argument of \\cite{luan2014} % can, in the isothermal approximation, also be reduced to a limit on the second moment of $n_e$, the emission measure: requiring the free-free optical depth $\\int\\md D\\,n_e^2\\bar\\alpha_\\mathrm{ff}(\\Theta)$, with $\\bar\\alpha_\\mathrm{ff}$ being the absorption coefficient of hydrogen plasma at unit density, to be below a reasonable value $\\kappa\\sim1$ implies \\bea\\label{luanlimit} a_2<\\frac{3^{5/2}c\\kappa}{2^{11/2}\\pi^{5/2}r_ef_0^2}\\frac{\\Theta^{3/2}}{\\tilde{g}_\\mathrm{ff}(\\Theta)} \\eea where $\\tilde{g}_\\mathrm{ff}\\sim10$ is a slowly-varying Gaunt factor, $\\Theta=k_BT/m_ec^2$ -- temperature in units of electron rest energy and $\\kappa$ the allowed optical depth.  The limit set by the allowed deviation $||\\delta\\hat t||$ from the low-density dispersion curve \\bea\\label{climit} a_2<\\frac{||\\delta\\hat t||}{\\max\\limits_x\\, (x-\\langle x\\rangle)^2} \\eea is independent of $\\Theta$ and $f_0$ but~(\\ref{luanlimit}) remains superior at temperatures up to $1.5\\times10^6\\,\\mathrm{K}$, assuming $\\kappa=1$, \\cite{thorntonetal2013} frequency coverage and sensitivity to  deviations at $||\\delta\\hat t||=3\\,\\mathrm{ms}$ level -- five times the resolution in Figure~\\ref{figure:dcurves}.  This might be a reason why no deviations from the standard dispersion curves have been reliably detected -- and ease some worries regarding FRB candidate selections. Concerns have been raised that selection criteria used might be effectively imposing cold plasma dispersion -- or indeed, extra-terrestrial origin -- on FRBs by only selecting candidates that conform to it from potentially a much broader population that does not. While some sort of selection is inevitable in the presence of noise and interference, the scope adopted might be too narrow for a new class of objects of as yet unknown origin, it has been argued. Superiority of~(\\ref{luanlimit}) over~(\\ref{climit}) implies that selection is not restrictive in relation to plasma density; if the signal emerges from the plasma unobsorbed, it would have the dispersion curve that does not deviate from the low-density limit at the resolution of current surveys. However, this will not necessarily be the case at lower frequencies or when the time resolution is improved. At $100\\,\\mathrm{MHz}$ the non-standard dispersion curve might show up for plasmas as cold as $T\\sim10^4\\,\\mathrm{K}$ (cf. note added in proof).  More speculatively, one might extend the reasoning above to the idea of FRB stellar origin itself. Could they originate in a relatively frequent phenomenon in dense stellar envelopes, of which FRBs is a fraction that happens at a density sufficiently low to let the FRB make it to the observer?\\footnote{\\cite{kulkarnietal2014} acknowledge, parenthetically, similar possibility  in relation to the absorption within FRB host galaxies.} If so, the estimates given above suggest the envelopes have to be rather hefty, extending to tens and hundreds solar radii. And although coronae and ionised outflows of these scales are not out of the question, they are certainly less numerous than the main sequence or flaring stars."
    },
    "1402/1402.2277_arXiv.txt": {
        "abstract": "Chandra  data  in  the COSMOS,  AEGIS-XD  and  4\\,Ms Chandra Deep Field South  are combined with multiwavelength photometry available in those  fields to determine the rest-frame  $U-V$ vs $V-J$ colours of X-ray AGN hosts in the redshift intervals $0.1<z<0.6$ (mean $\\overline{z}=0.40$) and $0.6<z<1.2$ (mean $\\overline{z}=0.85$).  This combination of colours provides an effective and least model-dependent means  of  separating  quiescent  from  star-forming,  including  dust reddened, galaxies.   Morphological information emphasises differences between AGN populations split by  their $U-V$ vs $V-J$ colours. AGN in quiescent  galaxies  consist   almost  exclusively  of  bulges,  while star-forming  hosts  are equally  split  between  early and  late-type hosts.  The  position of AGN  hosts on the  $U-V$ vs $V-J$  diagram is then  used to  set limits  on the  accretion density  of  the Universe associated with  evolved and star-forming systems  independent of dust induced biases.   It is found  that most of  the black hole  growth at $z\\approx0.40$  and  0.85   is  associated  with  star-forming  hosts. Nevertheless,  a  non-negligible  fraction  of  the  X-ray  luminosity density,  about  15-20\\%, at  both  $\\overline{z}=0.40$  and 0.85,  is taking place in galaxies in the quiescent region of the $U-V$ vs $V-J$ diagram.  For  the low-redshift  subsample, $0.1<z<0.6$, we  also find tentative evidence,  significant at the  $2\\sigma$ level, that AGN  split by their  $U-V$   and  $V-J$  colours  have   different  Eddington  ratio distributions.  AGN in blue  star-forming hosts dominate at relatively high Eddington ratios.  In contrast, AGN in red quiescent hosts become increasingly important  as a fraction  of the total  population toward low Eddington  ratios.  At higher redshift,  $z>0.6$, such differences are significant at the $2\\sigma$ level only for sources with Eddington ratios $\\ga10^{-3}$.   These findings are consistent  with scenarios in which  diverse accretion  modes are  responsible for  the  build-up of supermassive black holes at the centres of galaxies.  We compare these results with the predictions  of the {\\sc galform} semi-analytic model for  the  cosmological evolution  of  AGN  and  galaxies.  This  model postulates  two black  hole fuelling  modes,  the first  is linked  to star-formation events and the  second takes place in passive galaxies. {\\sc galform} predicts  that a substantial fraction of  the black hole growth  at $z<1$ is  associated with  quiescent galaxies,  in apparent conflict with the observations.  Relaxing the strong assumption of the model that passive AGN hosts have zero star-formation rate could bring those predictions in better agreement with the data. ",
        "introduction": "In recent years observational data established that supermassive black holes   (SMBHs)    are   nearly   ubiquitous    in   local   spheroids \\citep[e.g.][]{Magorrian1998}.  Moreover, correlations were discovered between the masses  of those black holes and  the stellar component of the  bulges  in  which  they  reside  \\citep[e.g.][and references therein]{Kormendy_Ho2013}.   These empirical correlations have been combined with large galaxy surveys to place tight constraints  on the mass function of  dormant SMBHs in the nearby   Universe  \\citep[e.g.][]{Kelly_Merloni2012}.    What  remains unclear however, is how the  relic SMBHs we observe in nearby galaxies grow  their  mass  across  cosmic  time.  One  way  to  approach  this question is to  conduct population studies of the  galaxies that host active   SMBHs   at  different   redshifts.    The  properties   (e.g. morphology,  environment)  of the  galaxies  with  an Active  Galactic Nucleus (AGN) provide information  on the physical conditions on large (kpc to Mpc) scales which may be relevant to the fuelling of the SMBH.  Morphological studies  for example, find  that X-ray AGN hosts  in the redshift range $z\\approx0.5-2$ have diverse morphologies (spiral, elliptical, disturbed) with a relative mix that is similar to that of mass-matched  non-AGN  galaxy samples  \\citep[e.g.][]{Georgakakis2009, Cisternas2011, Kocevski2012}.  This  suggests that major mergers, which  are expected  to be  associated with  morphologically disturbed systems, cannot  be the only channel  for growing black  holes at the centres of  galaxies.  Other  mechanisms, e.g.  minor  interactions or secular processes,  must also contribute  to the accretion  density of the  Universe.   This conclusion  is  also  supported  by large  scale structure  studies, which estimate  mean dark  matter halo  masses for X-ray AGN  in the range $\\rm  \\log (M/M_\\odot)\\approx12.5-13.5$.  This mass  interval is  larger than  expected  if black  hole accretion  is triggered   by   major   mergers   only   \\citep[e.g.][]{Allevato2011, Mountrichas2012,      Mountrichas2013}.       Recent      work      by \\cite{Fanidakis2013} indeed shows that  the clustering of X-ray AGN at $z\\la1.5$  is consistent  with two  channels for  growing  SMBHs.  The first is associated with star-formation  events in the host galaxy and the second is related to  quiescent galaxies in massive halos.  In the modelling of  \\cite{Fanidakis2013} star-formation  is a proxy  to cold gas availability.   Galaxies with abundant cold gas  supplies can form stars and grow their central black holes at a high rate.  In contrast, the SMBHs  of evolved galaxies  that are devoid  of cold gas  can only grow  slowly  via  the  accretion  of  hot  gas  from  a  quasi-static atmosphere.   Evidence   for  a   dichotomy  in  the   accretion  rate distribution  of  narrow  optical   emission-line  AGN  based  on  the star-formation  history of their  hosts is  reported at  low redshifts \\citep{Kauffmann_Heckman2009}.   This finding further  supports claims for diverse  AGN fuelling modes  and suggests that  the star-formation properties  of AGN hosts  hold important  information on  the physical conditions under which black holes at the centres of galaxies build-up their mass.  The   evidence  above   has  motivated   efforts  to   understand  the star-formation level  of AGN hosts  at higher redshift to  explore how black  holes are  fuelled as  a function  of cosmic  time.  Population studies have established that the build-up of black holes and galaxies are related in  a statistical sense when integrated  in a cosmological volume.  The star-formation rate density \\citep{Hopkins_A2006} and the accretion  luminosity  density  \\citep{Aird2010} follow  very  similar evolution  patterns with redshift.   There is  also evidence  that the cosmological  evolution of  the AGN  space density  is related  to the increase  with redshift  of the  average specific  star-formation rate (star-formation   rate   per    unit   stellar   mass)   of   galaxies \\cite[e.g.][]{Georgakakis2011}.     Far-IR/sub-mm   observations   with Herschel   extended  measurements  of   the  star-formation   rate  of individual AGN to high redshift and also bright accretion luminosities (e.g. luminous  QSOs) where  other indicators (e.g.   optical spectral features, broad-band colours) become  unreliable. Although it does not appear that there is a  one-to-one correspondence between the level of star-formation   and    the   accretion   power   \\citep{Mullaney2012, Rosario2012}, X-ray  AGN are on average associated  with galaxies on the  main  star-formation  sequence \\citep{Santini2012,  Mullaney2012, Rovilos2012, Rosario2013}.  In  a   typical  Herschel   far-IR/sub-mm  survey  field   however,  a substantial  fraction of  the  AGN population  lies  below the  formal detection limit.   Stacking the far-IR/sub-mm counts  at the positions of X-ray sources  is used extensively to reach  deeper flux limits and explore  the  star-formation  properties  of  the  entire  population. Although valuable, this approach  has the limitation that it estimates only the mean far-IR/sub-mm properties  of AGN hosts and provides only limited   information  on  the   underlying  distribution.    A  small sub-population of AGN not associated with high specific star-formation rate events  is likely  to be averaged  out in  far-IR/sub-mm stacking studies.  Optical observations for example, show that a large fraction of  the X-ray AGN  hosts at  $z\\la1$ lie  in the  red sequence  of the colour-magnitude  diagram \\citep[e.g.][]{Hickox2009, Georgakakis2011}, which includes  a large fraction  of passive galaxies.   Although dust can redden the  broad-band colours of galaxies, it  cannot account for the entire population  of AGN hosts on the red  sequence of the colour magnitude  diagram.    \\cite{Cardamone2010cmd}  for example,  find evidence for  a bi-modal $U-V$ rest-frame colour  distribution for AGN hosts at  $z\\approx1$, once  the impact of  dust is accounted  for via fitting  templates  to  the  observed spectral  energy  distributions. \\cite{Mignoli2004}  argue that  obscured X-ray  selected QSO  hosts at $z\\approx1-2$ have  rest-frame optical light profiles  that follow the de vaucouleurs  law.  This  is interepreted as  evidence that  a large fraction of the obscured QSO  population at $z\\approx1-2$ is hosted by bulge-dominated  galaxies, possibly  quiescent  ellipticals \\citep[but see ][]{Hutchings2002}.   In this paper we place limits on the fraction of the accretion density of the Universe associated with quiescent, low specific star-formation galaxies  in the  redshift range  0.1--1.2.   X-ray data  are used  to select AGN and their rest-frame  broad-band colours are adopted as the least  model-dependent  method  to  discriminate between  evolved  and actively star-forming  hosts.  Dust reddening issues  are mitigated by placing  X-ray AGN  hosts on  the $U-V$  vs $V-J$  (UVJ) colour-colour diagram \\citep{Williams2009, Patel2012}.   This combination of colours is  least sensitive  to  dust extinction  and  has been  shown to  be effective in  separating early-type, low-specific  star-formation rate galaxies from  actively star-forming, including  dust-reddened systems \\citep{Williams2009}. Throughout  this paper we adopt $\\rm  H_{0} = 70 \\,  km  \\,  s^{-1} \\,  Mpc^{-1}$,  $\\rm  \\Omega_{M}  = 0.3$  and  $\\rm \\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We  use  the  $UVJ$  diagram  to split  AGN  hosts  into  star-forming (including dust reddened) and quiescent. The host-galaxy morphology of the  two sub-populations  is found  to be  different,  suggesting that selection on the  $UVJ$ diagram provides a means  of identifying SMBHs that grow their mass under  distinct physical conditions. AGN hosts in the  quiescent  region  of  the  $UVJ$ diagram  are  early-type  bulge dominated  galaxies.  In  contrast, star-forming  AGN hosts  include a large fraction  (about 50\\%) of  late-type systems.  We  also estimate the  fraction  of the  accretion  density  associated  with those  two classes of  hosts at  redshifts $z\\approx0.40$ and  0.85. Most  of the accretion density  at those redshifts is taking  place in star-forming hosts. Nevertheless,  about 15-20\\% of  the AGN luminosity  density is associated with  galaxies in  the quiescent part  of the  UVJ diagram. There is also  evidence, significant at the $2\\sigma$  level, that AGN in the low redshift  subsample ($0.1<z<0.6$) have accretion properties that depend  on the  level of star-formation  of their hosts.   AGN in star-forming hosts  dominate at high Eddington ratios,  while those in quiescent  hosts become  increasingly important  toward  low Eddington ratios.  At higher redshift, $0.6<z<1.0$, such differences are present at  the $2\\sigma$  level  only  for AGN  with  Eddington ratios  $\\log \\lambda_{Edd}\\ga-3$.  These results are  consistent with two modes for growing  black  holes  that  take  place in  galaxies  with  different star-formation  properties.  We  compare those  observations  with the predictions  of {\\sc  galform} SAM,  which postulates  two  black hole growth channels, one linked  to star-formation and the other occurring in passive systems.  This SAM predicts a fraction of accretion density in quiescent  hosts that  is larger than  the observed.  We  also find that the  evolution with redshift  of the X-ray luminosity  density of hot-halo/starburst mode AGN is inconsistent with the observations.  We argue that  these discrepancies could be attributed  to the assumption of {\\sc galform}  that the hot-halo accretion mode  is not accompanied with some level  of star-formation in the host  galaxy.  Relaxing this requirement could bring the model in better agreement with the data."
    },
    "1402/1402.3271_arXiv.txt": {
        "abstract": "We sift the impact of the recent Higgs precise measurements, and recent dark matter direct detection results, on the dark sector of an electroweak extension of the Standard Model that has a complex scalar as dark matter. We find that in this model the Higgs decays with a large branching ratio into dark matter particles, and charged scalars when these are kinematically available, for any coupling strength differently from the so called Higgs portal. Moreover, we compute the abundance and spin-independent WIMP-nucleon scattering cross section, which are driven by the Higgs and $Z^{\\prime}$ boson processes. We decisively exclude the $1-500$~GeV dark matter window and find the most stringent lower bound in the literature on the scale of symmetry breaking of the model namely $10$~TeV, after applying the LUX-2013 limit. Interestingly, the projected XENON1T constraint will be able to rule out the entire $1$~GeV-$1000$~GeV dark matter mass range. Lastly, for completeness, we compute the charged scalar production cross section at the LHC and comment on the possibility of detection at current and future LHC runnings. ",
        "introduction": "\\label{introduction}  The nature of the dark matter (DM) is one of the greatest puzzles in current science, once the DM constitutes approximately $23\\%$ of the Universe budget. There are promising ongoing searches aimed to detect and find the nature of the DM that permeates the Universe. There are many dark matter candidates in the literature, but the most seemingly promising ones are the so called WIMPs (Weakly Interacting Massive Particles) for having a thermal cross section at the electroweak scale, naturally addressing the structure formation process, and being predicted in many interesting particle physics models.  There are four different methods to infer the presence or detect theses WIMPs known as indirect detection, direct detection, colliders and cosmological observations. Indirect detection searches have found some excess events in the gamma-ray emission \\cite{gammaray} and in the cosmic ray emission \\cite{cosmicray} which might be explained by annihilation of WIMPs in our galaxy \\cite{IDexplanations}, which is in contradiction with recent dwarf galaxies constraints \\cite{dwarfsAlma}. Likewise, some direct detection experiments such CoGeNT \\cite{cogent}, DAMA \\cite{dama}, CRESST \\cite{cresst} and most recently CDMSII-Si \\cite{cdms} have observed some excess events consistent with WIMP scatterings \\cite{ddwimps}. Due to some possible leakage of background events into the signal region at low energies and the non-observation of such events in the XENON \\cite{xenon} and LUX \\cite{lux} experiments, those events do not constitute an irrefutable DM signal \\cite{backgrounds}. Furthermore, there are cosmological measurements of the Cosmic Microwave Background that revealed some degree of dark radiation observed in the Planck data \\cite{planck}, among other satellites, \\cite{satelites} that may constitute an evidence for a sub-dominant non-thermal production of DM \\cite{nonthermal}. Lastly, collider data, which provide an important and complementary method to infer the nature of the dark matter have not observed any positive signal for a stable particle and just bounds on the mass and coupling strengths had been derived \\cite{Willshepherd}.  In this work we will focus on a compelling extension of the Standard Model (SM) namely 331LHN, that might address these evidences. 331LHN stands for a electroweak extension of the SM where doublets are replaced by triplets, both in the scalar and the fermion sector. This proposal has been able to endure all electroweak precise measurements and reproduce the SM results concerning the Higgs signal strength \\cite{331DM2} as oppose to other 331 model extensions which predict a $H\\rightarrow \\gamma \\gamma$ enhancement \\cite{hgamma}. It also has a rich particle spectrum comprised of charged scalars , gauge bosons, sterile neutrinos and exotic quarks, with interesting phenomenological aspects, which had been investigated elsewhere \\cite{gaugebosonsLHC}. Furthermore, this model does have a plausible DM candidate able to explain the gamma-ray excess observed in the Fermi-LAT data at the Galactic Center \\cite{gammaray} differently to other versions \\cite{331minimal,331econo,331susyecono,331two,331susyecono2,331g2muon,Caetano:2012qc,Ferreira:2013nla,Dias:2013kma,Dong:2014wsa}; and offers a plausible mechanism to account for the dark radiation observed by the Planck Collaboration, through a sub-dominant non-thermal production of WIMPs,  while evading structure formation, Big Bang Nucleosynthesis and CMB bounds among others \\cite{331darkradiation}. An extensive analysis concerning the heavy fermions present in the model has been done, and stringent bounds on the mass of the lightest sterile neutrino have been found as a function of the $Z^{\\prime}$ mass in Ref.\\cite{331DM3}, and in a model independent fashion in Ref.\\cite{DarkZp}. It is important to stress that such constraints on the $Z^{\\prime}$ mass {\\it do apply, at some level, to all 331 models}, that have fermions as DM candidates, as discussed in Fig.7 of Ref.\\cite{331DM3}, and are complementary to others coming from colliders \\cite{331collider,331collider2}, FCNC \\cite{331FCNC}, muon decay \\cite{331muon}, top decay \\cite{331top} analyses, and oblique STU parameters \\cite{Liu:1993fwa,Martinez:2009ik}.  That being said, here we will discuss the 331LHN model % which has two not co-existing DM candidates \\cite{331DM2,331DM1}. Our purpose is to derive constraints on the dark sector of this well motivated model in the light of the present bounds in the Higgs signal strength and DM observables. In particular, we will exclude DM masses below $500$~GeV, and discuss the impact of this exclusion on the dark sector of the model with focus on the charged scalar which predicted in this model.  The paper is organized as follows: in section \\ref{sec1} we briefly introduce the 331LHN model. In section \\ref{darksector} we derive bounds on the dark sector of the model, and in section \\ref{LHC} we comment about the possibility of detection at current and future LHC runnings. Finally we present our conclusions in section \\ref{sec:conclusion} ",
        "conclusions": "\\label{sec:conclusion}  We have examined bounds on the dark sector of a 331 model known as 331LHN, which contains heavy neutral fermions ($N_{a}$) and a complex scalar dark matter particle in its spectrum, based on the current Higgs and direct dark matter detection data. The model is comprised of three scalar triplets, and interestingly, all of them couple to the Higgs boson. Therefore, we found a lower bound on the mass of these scalars by imposing the LHC constraints concerning the Higgs signal strength. In particular, we found that it requires the mass of the WIMP ($\\phi$) and the charged scalars $(h_1^{\\pm})$ to be all heavier than $M_H/2$~GeV, regardless of the coupling values used, differently from the so called Higgs portal.  We have also computed numerically the abundance and scattering cross section of the WIMP taking into account all possible amplitudes. Combining the Higgs and DM constraints we found the most stringent constraints in the literature on this model. Our main results read:  (i) Close DM masses below 500GeV;\\\\ (ii) Exclude light WIMPs as an explanation for the Galactic Center excess;\\\\ (iii) Find a lower bound of 10 TeV on the symmetry breaking scale;   Moreover, the projected XENON1T bounds are expected to fiercely rule out the entire 1GeV-1TeV dark matter region. Therefore, combining the Higgs and dark matter data, we decisively close the light dark matter window in this model and showed that the scale of symmetry breaking of this model has to live at the $\\sim 10$~TeV in order to have a viable DM candidate.  Lastly, for completeness we have computed the production cross section of the charged scalars $h_1^{\\pm}$ at the LHC, which is driven by the $Z^{\\prime}$ mass, and concluded that these charged scalars might be within reach of the LHC at 14~TeV as shown in FIGS.\\ref{fig3} and \\ref{fig4}."
    },
    "1402/1402.7102_arXiv.txt": {
        "abstract": "The gas dynamics of protoplanetary disks (PPDs) is largely controlled by non-ideal magnetohydrodynamic (MHD) effects including Ohmic resistivity, the Hall effect and ambipolar diffusion. Among these the role of the Hall effect is the least explored and most poorly understood. In this series, we have included, for the first time, all three non-ideal MHD effects in a self-consistent manner to investigate the role of the Hall effect on PPD gas dynamics using local shearing-box simulations. In this first paper, we focus on the inner region of PPDs, where previous studies (Bai \\& Stone 2013, Bai, 2013) excluding the Hall effect have revealed that the inner disk up to $\\sim10$ AU is largely laminar, with accretion driven by a magnetocentrifugal wind. We confirm this basic picture and show that the Hall effect introduces modest modifications to the wind solutions, depending on the polarity of the large-scale poloidal magnetic field ${\\mb B}_0$ threading the disk. When ${\\mb B}_0\\cdot{\\mb\\Omega}>0$, the horizontal magnetic field is strongly amplified toward the disk interior, leading to a stronger disk wind (by $\\sim50\\%$ or less in terms of the wind-driven accretion rate). The enhanced horizontal field also leads to much stronger large-scale Maxwell stress (magnetic braking) that contributes to a considerable fraction of the wind-driven accretion rate. When ${\\mb B}_0\\cdot{\\mb\\Omega}<0$, horizontal magnetic field is reduced, leading to a weaker disk wind (by $\\lesssim20\\%$) and negligible magnetic braking. Moreover, we find that when ${\\mb B}_0\\cdot{\\mb\\Omega}>0$, the laminar region extends farther to $\\sim15$ AU before the magneto-rotational instability sets in, while for ${\\mb B}_0\\cdot{\\mb\\Omega}<0$, the laminar region extends only to $\\sim3$ AU for a typical accretion rate of $\\sim10^{-8}-10^{-7}M_{\\bigodot}$ yr$^{-1}$. Scaling relations for the wind properties, especially the wind-driven accretion rate, are provided for aligned and anti-aligned field geometries. Issues with the symmetry of the wind solutions and grain abundance are also discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5870_arXiv.txt": {
        "abstract": "The search for and potential identification of dark matter (DM) will require a simultaneous, multi-pronged approach with important roles to be played by the LHC, both direct and indirect DM detection experiments and, eventually, the ILC. The 19-parameter p(henomenological)MSSM can provide a broad framework for complementarity studies of neutralino DM. In this talk, I summarized the sensitivity of these searches at the 7, 8 and 14 TeV LHC, combined with those from \\Fermi, CTA, IceCube/DeepCore, COUPP, LZ and XENON. The strengths and weaknesses of each of these searches is briefly discussed as is their interdependent roles in probing the pMSSM model parameter space. I also comment on the future role of the ILC in exploring the detailed nature of neutralino DM. ",
        "introduction": "Although the reality of DM is no longer in question, its true nature remains as mysterious as ever. Because of a numerical `miracle', the simplest and most often studied possibility is that DM is a Weakly Interacting Massive Particle (WIMP) whose thermal relic abundance is set by its annihilation cross section in the early universe. Beyond simplicity, one of the reasons WIMPs are so attractive is that they can be sought using multiple experimental techniques: through direct production at accelerators, direct detection (DD) by scattering off nuclei in experiments underground and indirect detection (ID) of DM annihilation products coming from the sky. Even if convincing DM signals are discovered in one of these channels in the future, however, it is likely that complementary information from all three of these sources will be necessary to elucidate the detailed properties of DM. So far, no such conclusive signals for DM have yet been observed and the parameter space for WIMPs continues to slowly shrink though much yet remains to be explored.  In order to make use of possible complementary information on the properties of DM one must have a framework that completely describes how the DM interacts with the various particles of the Standard Model (SM). Here there are two different approaches: in a bottom up, `model-independent' framework, one posits a simplified model with only the minimal particle content, \\ie, the DM particle itself as well as a mediator connecting the DM with the SM.  Furthermore, provided the mediator is itself sufficiently massive to be integrated out, the resulting interactions are described by a set of higher dimension operators, \\ie, an EFT. In the situation where the EFT is applicable this approach has been shown to be quite powerful, especially for DD and ID searches. However, in the case of colliders such as the LHC this EFT approach frequently breaks down and one needs to revert back to the original simplified model for consistency. Of course if DM is not the only light(ish) state even this simplified model approach will itself fail. A second possibility is to adopt a UV-complete scenario where the DM particle is accompanied by a potentially wide range of other new states of various masses all of which can interact with the SM through electroweak and/or strong interactions. Unlike in the EFT or simplified model approaches, in this case we are likely forced to buy into a rather complex scenario with a multi-dimensional parameter space that needs to be examined.  Of the possible UV-complete scenarios predicting WIMPs, R-parity conserving supersymmetry (SUSY) appears to be the most attractive with the MSSM being its simplest realization. Here the lightest stable SUSY particle (LSP) that is both electrically neutral and a color-singlet, often the lightest neutralino ($\\chi_1^0$) as we will assume in the discussion below, can be identified as a possible WIMP DM candidate. While DM searches are obviously focused on the properties of the LSP itself, within this UV framework the properties of many if not all of the other superparticles as well as those of the members of the extended SUSY Higgs sector can also play important roles in the various searches. Thus it is impossible to completely separate DM searches from searches for and the examination of the rest of the SUSY spectrum. Furthermore, even in the simplest SUSY scenario, the MSSM, the number of free parameters ($\\sim$ 100) is simply too large to perform a study in all generality so we need to restrict our attention in some way without (hopefully) losing contact with any of the relevant physics.  To at least partially circumvent this problem, in the analysis discussed here we restrict ourselves to the 19-parameter p(henomenological)MSSM as is fully discussed detail elsewhere{\\cite {us}}. For completeness, the 19 pMSSM parameters and their ranges employed in the results below are presented in Table~\\ref{ScanRanges}. To study the pMSSM, we essentially throw darts into this large parameter space, generating many millions of random model points, with each point corresponding to a specific set of values for these 19 parameters. These individual `models' are then subjected to a large set of collider, flavor, precision measurement, dark matter and theoretical constraints with roughly 225k models surviving these for further study here. One interesting result of employing these scan ranges is that the LSP is frequently (but not always) close to being an electroweak eigenstate, \\ie, a relatively pure bino, wino or Higgsino, which has important consequences for DM searches.  Below, we will briefly overview the sensitivities of both present and future experiments which employ different approaches in searching for SUSY dark matter: collider production, as well as direct and indirect detection, to see how well this pMSSM parameter space is `covered'. The analyses will {\\it not} assume that the thermal relic density as calculated for the LSP necessarily saturates the WMAP/Planck value, $\\Omega h^2\\simeq 0.12$, to allow for the possibility of multi-component DM. For example, the axion introduced to solve the strong CP problem might make up a substantial amount of the DM. In the pMSSM the nature of the neutralino LSP can vary widely as one moves about in this large parameter space. Even if the LSP is is (in)directly observed at the LHC it is likely to be difficult to fully determine its electroweak properties even with the supplementary information provided by the DD and/or ID DM experiments. These properties are of course critical to identifying the LSP as a true component of DM. It is likely that this information can only be provided by the ILC, which can fully study the nature of the LSP provided it is kinematically accessible, in order to complete a consistent picture of DM.   \\begin{table} \\centering \\begin{tabular}{|c|c|} \\hline\\hline $m_{\\tilde L(e)_{1,2,3}}$ & $100 \\gev - 4 \\tev$ \\\\ $m_{\\tilde Q(q)_{1,2}}$ & $400 \\gev - 4 \\tev$ \\\\ $m_{\\tilde Q(q)_{3}}$ &  $200 \\gev - 4 \\tev$ \\\\ $|M_1|$ & $50 \\gev - 4 \\tev$ \\\\ $|M_2|$ & $100 \\gev - 4 \\tev$ \\\\ $|\\mu|$ & $100 \\gev - 4 \\tev$ \\\\ $M_3$ & $400 \\gev - 4 \\tev$ \\\\ $|A_{t,b,\\tau}|$ & $0 \\gev - 4 \\tev$ \\\\ $M_A$ & $100 \\gev - 4 \\tev$ \\\\ $\\tan \\beta$ & 1 - 60 \\\\ \\hline\\hline \\end{tabular} \\caption{Scan ranges for the 19 parameters of the pMSSM with a neutralino LSP employed in the present analysis. All parameters are scanned with flat priors~\\cite{us}.} \\label{ScanRanges} \\end{table}    \\begin{figure}[htbp] \\centerline{\\includegraphics[width=7.0in]{mLSP_Omega_None_Eigenstates.png}} \\vspace*{-0.10cm} \\caption{Thermal relic density as a function of the LSP mass in our pMSSM model set, as generated, color-coded by the electroweak properties of the LSP as discussed in the text.} \\label{fig00} \\end{figure}    Figure~\\ref{fig00} provides an overview of the properties of the many LSPs in our model set. Here we see the thermal relic densities of the LSPs as a function of their mass with the color-coding reflecting their electroweak eigenstate content. There are many things to note here that will be important for later consideration. Essentially every possible mechanism to obtain (or to lie below) the WMAP/Planck relic density is present here: ($i$) The set of models with low LSP masses (forming `columns' on the left-hand side of the Figure) correspond to bino-Higgsino admixtures which achieve a sufficiently low relic density by resonant annihilation through the $Z,h$-funnels; these sometimes appear as pure binos if the Higgsino fraction is very small{\\footnote {Here, `pure' means having an eigenstate fraction $\\geq 90\\%$. Points shown as bino-wino, bino-Higgsino, or wino-Higgsino mixtures have less than $2\\%$ Higgsino, wino, or bino fraction, respectively. Mixed points have no more than $90\\%$ and no less than $2\\%$ of each component.}}. ($ii$) The bino-Higgsino LSPs saturating the relic density in the upper left region of the figure are of the so-called `well-tempered' variety. ($iii$) The pure bino models in the upper middle region of the Figure are bino co-annihilators (mostly but not exclusively with with sleptons) or those which annihilate resonantly through the $A$-funnel. ($iv$) The green (blue) bands are pure Higgsino (wino) models that saturate the relic density bound (using perturbative calculations which do not include Sommerfeld enhancement) near $\\sim 1(1.7)$ TeV and have very low relic densities for lighter LSP masses. Wino-Higgsino hybrids are seen to lie between these two cases as expected. ($v$) A smattering of models with additional (or possibly multiple) aids to their annihilation are loosely distributed in the lower right-hand corner of the Figure. As we will see, many of the searches for DM are particularly sensitive to one or more of these LSP categories. We now turn to a brief discussion of the individual searches. ",
        "conclusions": "These results lead to a number of interesting conclusions:    \\begin{itemize}   \\item{Even if the LSP {\\it does not} make up all of the DM, it can still be observed in both direct and indirect detection experiments as well as neutrino experiments such as IceCube. Of course, searches at the LHC are not influenced by the LSP relic density.}  \\item{The set of models remaining after all the searches are performed, yielding null results,  that saturate the thermal relic density consist almost uniquely of those with (co)annihilating bino LSPs.}  \\item{SI direct detection via LZ, CTA, and the LHC do most of the heavy lifting in terms of complementary searches covering the pMSSM parameter space.}  \\item{Multiple/overlapping searches allow for extensive parameter space coverage which will be of particular importance if a DM signal is observed.}  \\item{Most of the experiments are seen to provide complementary probes of the pMSSM parameter space.}  \\item{The strength of the LHC component in these searches increases significantly with the inclusion of the 14 TeV LHC searches. However, these searches dependent on the rest of the model spectrum and therefore do not provide complete coverage of any given LSP scenario, in contrast to DM searches which rely more directly on the LSP properties.}  \\item{The ILC has an important role to play in determining the electroweak properties of the the LSP after discovery to complete the whole DM picture.}  \\end{itemize}   In summary, the pMSSM provides an excellent tool for studying complementarity between different approaches to the search for and exploration of dark matter. Hopefully, DM will soon be discovered so that we can employ the complementary probes discussed above to ascertain its true nature."
    },
    "1402/1402.2457_arXiv.txt": {
        "abstract": "Observing the Sun at high time and spatial scales is a step towards understanding the finest and fundamental scales of heating events in the solar corona. The Hi-C instrument has provided the highest spatial and temporal resolution images of the solar corona in the EUV wavelength range to date. Hi-C observed an active region on 11 July 2012, which exhibits several interesting features in the EUV line at 193\\AA: one of them is the existence of short, small brightenings ``sparkling\" at the edge of the active region; we call these EUV Bright Dots (EBDs). Individual EBDs have a characteristic duration of 25s with a characteristic length of 680 km. These brightenings are not fully resolved by the SDO/AIA instrument at the same wavelength, however, they can be identified with respect to the Hi-C location of the EBDs. In addition, EBDs are seen in other chromospheric/coronal channels of SDO/AIA suggesting a temperature between 0.5 and 1.5 MK. Based on their frequency in the Hi-C time series, we define four different categories of EBDs: single peak, double peak, long duration, and bursty EBDs. Based on a potential field extrapolation from an SDO/HMI magnetogram, the EBDs appear at the footpoints of large-scale trans-equatorial coronal loops. The Hi-C observations provide the first evidence of small-scale EUV heating events at the base of these coronal loops, which have a free magnetic energy of the order of 10$^{26}$ erg. ",
        "introduction": "A long standing problem in solar physics is how the solar corona is heated up to a temperature of few million degrees while the photosphere is at about 6000 K. Several models have been developed either to heat the coronal plasma locally or to transport heat from the lower layers of the solar atmosphere into the corona (see for instance review by \\citeauthor{klim06} \\citeyear{klim06}), however, these models often lack an explanation for an average temperature of 1MK for the entire corona and its sustainability during a solar cycle. This strongly suggests that several mechanisms may be at play, each of them providing a substantial amount of heat to the corona and, which can vary during the activity cycle of the Sun. Those models also invoke smaller spatial and time scales than the ones provided by the current observations.  So, the critical issue is really to answer the questions: where is the source of heat located? and where and how does the heating take place? Along a given magnetic structure (often a coronal loop), three scenarios are often considered: the release of energy is located (i) at the footpoints of the structure in the photospheric layer which is the most favorable location owed to the importance of the photospheric motions and the source of material and energy transported from the convection zone, (ii) at the apex of the structure, and (iii) uniformly along the structure. Those three scenarios imply that the heat is transported along the magnetic structures, and that we actually observe loop-like structures in the corona. In terms of observations, it has been difficult to disentangle the actual location of the heat source. As an example, several authors using a dataset obtained in soft X-rays by {\\em Yohkoh}/SXT have reached an opposite conclusion depending on how they dealt with the data reduction \\citep{mac00,pri00,reale02}. To determine the possible heating scenario, the authors looked at the profile of temperature along the coronal loop and compared with a theoretical model. These indirect methods were needed because of the coarse spatial and time resolutions of {\\em Yohkoh}/SXT (compared to what can be achieved nowadays). In this paper, with the development of new high spatial and time resolution instrumentation, we are looking for a direct evidence of the source of heat deposition, observationally charaterised by a small, short increase in intensity.  Theoretically, the different heating locations (apex, uniform or footpoint) arise from the idea that magnetic energy stored in the chromosphere/corona produces small, short bursts converting the magnetic energy into thermal and/or kinetic energy (rapidly converted into heat). This is the basic principle of \\citeauthor{par88}'s theory (\\citeyear{par88}) for heating the solar corona by nanoflares. A large number of those small events is needed to sustain and uniformly distribute the heat in the entire solar corona. The mechanism responsible for the release of magnetic energy could be, for instance, magnetic reconnection, wave mode coupling or turbulence. As a large amount of the free energy contained in the magnetic field structure of active regions is located at the base of the corona or in the chromosphere \\citep{reg07}, the most favorable heating location is at the footpoints of coronal loops. For the quiet Sun, the footpoints of the coronal strucutres are the most effected by the photospheric motions such as granular motions which can easily produce tangled/braided/twisted magnetic flux bundles. In this paper, we combine EUV observations and a magnetic field model to determine if Parker's model is a viable scenario, even if we cannot identify the mechanism(s) responsible for the energy release.  High-resolution ground-based instruments in the visible-wavelength domain allowed to observe new structures and phenomena in the lower layers of the solar atmosphere, especially small-scale brightenings or events such as spicules or Ellerman bombs \\cite[see e.g.,][]{sch03,ber04,rut04,rou05,bec07}. In conjunction, the investigation of the coronal plasma at EUV wavelengths is in constant improvement in terms of spatial and temporal resolution. For instance, the Atmospheric Imager Assembly (AIA; \\citeauthor{lemen12} \\citeyear{lemen12}) on board the Solar Dynamics Observatory (SDO; \\citeauthor{pes12} \\citeyear{pes12}) is largely contributing to a new era of fast (12s) small-scale (pixel size of 0.6\\arcsec) observations of the solar corona always discovering new phenomena. As a new step forward, the Hi-C instrument has provided high-temporal (5.5s) and high-spatial (pixel size of 0.1\\arcsec) resolution observations of the EUV corona. Using this dataset, \\citet{cirt13} have recently shown the existence of small-scale braided structures within the active region, and evidenced the magnetic reconnection process occurring within the braided structures.   In addition, two recent articles relying on Hi-C data have provided new insights in the existence of small-scale heating events. \\cite{tes13} have found that the active region moss is highly dynamic with a variability on timescales of 15s, and the moss brightenings lead to an energy release of 10$^{23}$ erg (order of magnitude consistent with the nanoflare model). Similarly, \\cite{win13} have investigated transient brightenings along small-scale transition-region loops. It was found that the loops have a diameter ranging from 0.9\\arcsec\\ to 1.1\\arcsec\\ with an average duration of 64s. \\cite{win13} have estimated that these brightenings  radiate an energy of about 10$^{24}$--10$^{25}$ erg again an energy estimate consistent with the nanoflare model.  In this paper, we report on the existence of an additional heat source for the coronal plasma: EUV, short-lived (25 s), small-scale ($<$1\\arcsec) brightenings taking place in the low layers of the solar atmosphere (log T = 5.5--6.5). We call these brightenings EUV bright dots (EBDs). This new discovery is made possible by the high spatial resolution and fast time cadence of the Hi-C data.  In Section~\\ref{sec:obs}, after introducing the Hi-C instrument, we characterize  the location, size, and thermal properties of EBDs as derived from the Hi-C channel and six SDO/AIA channels. We thus study their time evolution in the observed region (Section~\\ref{sec:dyn}), as well as their magnetic properties (Section~\\ref{sec:mag}) relying on a potential field extrapolation of one SDO/HMI magnetogram. In Section~\\ref{sec:em_loci}, the possible temperature of EBDs is investigated using the EM loci method. The conclusions are drawn in Section~\\ref{sec:concl}.  \\begin{figure} \\begin{center} \\includegraphics[width=1.\\linewidth]{fig1_small.eps} \\caption{Location of the region of interest containing the EBDs in the Hi-C FOV recorded at 18:53:32 UT on 11 July 2012.} \\label{fig:hic_context} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:concl}  The high resolution images obtained in the EUV by Hi-C allow us to determine the smallest brightenings observed to date in an active region in extreme ultra-violet light, called EUV bright dots (EBDs). These EBDs sparkle at the footpoints of large-scale coronal loops. We have grouped the EBDs into four categories depending on their recurrence. Except for the fourth category (successive, short-lived EBDs), individual EBDs which can be extracted from the Hi-C time series have the same characteristic duration (about 25s) and length scale (0.68 Mm). Using a potential field extrapolation of the SDO/HMI line-of-sight photospheric magnetic field, we have determined that the EBDs observed in the Hi-C FOV are located at the footpoints of large-scale magnetic loops connecting the active regions 11520 to a decaying active region in the Northern hemisphere. The above properties suggest that EBDs might be a signature of an impulsive release of energy located at the base of coronal loops.  In terms of energetics, the magnetic energy estimate is of the order of 10$^{26}$ erg which is 100 times more than required by the nanoflare model suggested by \\citet{par88}. However, our estimate is an upper bound for the free magnetic energy than can be released, thus suggesting that smaller scale structures can exist. Nevertheless, \\citet{tes13} and \\citet{win13} have estimated the radiated energy to 10$^{24}$ erg, which is compatible with the order of magnitude of the magnetic energy: according to \\citet{ems12} who performed a complete study of the energy partition in eruptive events, the magnetic energy budget can be two orders of magnitude larger than the radiated energy.  Even if it is not possible to compare quantitatively due to the lack of spectroscopic data overlying the Hi-C FOV, it is important to mention the observations of \\citet{har09}, which showed that small-scale brightenings at the footpoints of coronal loops can be observed with a strong redshift velocity close to the local sound speed. \\citet{har09} suggested that spatially unresolved ($<$720 km) structures could exist at the footpoint of coronal loops.  According to \\cite{walsh95}, the conduction time-scale can be expressed as $\\tau_c = pL^2 / \\kappa_0T^{7/2}$ where $p$ is the plasma pressure, $L$ the characteristic length, $\\kappa_0$ the thermal conductivity in the corona ($\\kappa_0 = 10^{-11}$), and $T$ is the temperature. The characteristic length $L$ of the loop is chosen to be the loop length along which the heating is assumed to take place: $L$ = 5 Mm. For a plasma pressure of 0.1 dyne$\\cdot$cm$^{-2}$ and a temperature range from 0.5 to 1.5 MK, the conductive time-scale is between 6 s at 1.5 MK and 5 min at 0.5 MK. A characteristic time of 25 s gives a characteristic temperature of 1 MK. Thus, the observed life-time of EBDs maybe compatible with the properties of a coronal plasma.  With the restricted observations of the Hi-C instrument (short flight and limited FOV), questions arise for further study of these elementary events: how often/common are EBDs in active regions? Are they always located at the footpoints of long coronal loops? What is their contribution towards the global heating of the corona? Can EBDs be associated with Type 2 spicules or even Ellerman bombs, both having similar time scales? Several of these questions can be addressed by finding criteria which will enable the detection of EBDs in SDO/AIA images going forward, and by using high resolution chromospheric spectroscopic observations from IRIS."
    },
    "1402/1402.5041_arXiv.txt": {
        "abstract": "{} { We are presenting here a follow-up study of the molecular gas and dust  in the environs of the star forming region \\ngc. This study aims at dealing with the interaction of the \\hii\\ region \\ngc\\  with its parental molecular cloud, and also with the star formation in the region, that was possibly triggered by the expansion of the ionization front against the parental cloud.    } {To analyze the molecular gas we use \\cob, \\coc, \\cod, and \\hcn\\ line data obtained with the on-the-fly technique from the APEX telescope. To study the distribution of the dust, we make use of unpublished images  at 870 $\\mu$m from the ATLASGAL survey and IRAC-GLIMPSE archival images. We use public 2MASS and WISE data to search for infrared candidate YSOs in the region. } {The new APEX observations allowed the substructure of  the molecular gas in the velocity range from $\\sim$ $-$28 to $-$23 \\kms\\ to be  imaged in detail. The morphology of the molecular gas close to  the nebula, the location of the PDR, and the  shape of radio continuum emission suggest that the ionized gas is expanding against its parental cloud, and confirm the ``champagne flow'' scenario. We have identified several molecular clumps and determined some of their physical and dynamical properties  such as density, excitation temperature, mass, and line width. Clumps adjacent to the ionization front are expected to be  affected by the \\hii\\ region, unlike those that are distant to it. We have compared the physical properties of the two kind of clumps  to investigate how the molecular gas has been affected by the \\hii\\ region. Clumps adjacent to the ionization fronts of \\ngc\\ and/or   the bright rimmed cloud   \\sfo\\ have been heated and compressed by the ionized gas, but their line width is not different to those that are too distant to the ionization fronts.  We identified several candidate  YSOs in the region. Their spatial distribution suggests that stellar formation might have been boosted by the expansion of the nebula. We discard the ``collect and collapse'' scenario and propose  alternative mechanisms such as radiatively driven implosion on pre-existing molecular clumps or small-scale Jeans gravitational instabilities. } {} ",
        "introduction": "It is accepted that OB associations have an enormous impact on the state of their environs. The interstellar medium (ISM) surrounding OB stars is expected to be strongly modified and disturbed  by their intense ultraviolet (UV) radiation field ($h\\nu>$ 13.6 eV). UV photons ionize the surrounding gas creating \\hii\\ regions and dissociate the molecular gas originating photodissociation regions (PDRs) \\citep{ht97}. Further, the surrounding neutral gas (either atomic or molecular), is compressed by the expansion of the \\hii\\ region and/or the action of stellar winds.  The compression of the cloud  could enhance the stellar formation via ``radiative driven implosion'' process (RDI;  \\citealt{lela94}) or even trigger it via ``collect and collapse'' process \\citep{elm77}. Therefore, when massive stars form inside a molecular cloud it is expected that they dominate the state of the parental cloud and consequently the further stellar formation process within. Indeed, it has been shown that a large fraction of stars originates at the peripheries of \\hii regions \\citep{pom09,ro09,ca09,deh10,va12,deh12}. Having this in mind, it is instructive to study the molecular gas adjacent to \\hii regions  since it  can provide important information on the interaction  between massive stars and their natal environments. Furthermore, a comparison between regions of the molecular cloud adjacent to an \\hii region with other sites away from it may provide significant understanding into the physical properties of the molecular gas that may impact the formation of stars.  \\begin{figure*} \\centering \\includegraphics[width=470pt]{ukst2.eps} \\caption{UKST red plate image of the nebula \\ngc\\ in  the area covered by the APEX data. The positions of \\ngc, the bright rimmed cloud \\sfo, and the open cluster Pis 17 are indicated.} \\label{fig:ukst} \\end{figure*}  The optical emission nebula NGC\\,3503 (= Hf 44 = BBW 335) is a small \\hii region located at  \\radec\\ = (11$^h$01$^m$16$^s$, $-$59\\gra50\\arcmin39\\arcsec) \\citep{dr88}, and placed at a distance of \\hbox{2.9 $\\pm$ 0.4 kpc} \\citep{pco10}. \\ngc\\  is  ionized by early B-type stars belonging to the open cluster  Pis 17 \\citep{H75,pco10,dvcca12} and is believed to be related to the bright-rimmed cloud (BRC)  {\\rm SFO}\\,62 \\citep{so91,Y99,T04,U09}, although the ionizing star/s of the BRC have not  been certainly identified.  In a recent work (\\citealt{dvcca12}; hereafter Paper I) we carried out a  multifrequency analysis in the environs of \\ngc. We analyzed the properties of the molecular gas using NANTEN $^{12}$CO (J=1$\\rightarrow$0)   (HPBW = 2.$'$7) observations, whilst the ionized gas was studied making use of radio continuum observations at 4800  and 8640 MHz,  with synthesized beams of \\hbox{23\\farcs 55 $\\times$ 18\\farcs 62}   and  \\hbox{14\\farcs 73 $\\times$ 11\\farcs 74}, respectively, carried out with  ATCA. The molecular line observations  revealed a molecular gas component of \\hbox{7.6 $\\times$ 10$^3$ \\msun}in mass having a mean radial velocity\\footnote{Radial velocities in this paper are always referred to the local standard of rest ({\\rm LSR})} of $-$24.7 \\kms (in agreement with the velocity of the ionized gas of \\ngc; \\citealt{g00}),  that  is associated with the nebula and its surroundings.  We reported an overdensity   centered at   \\radec\\ = (11$^h$01$^m$02.48$^s$, $-$59\\gra50\\arcmin04.6\\arcsec\\  \\hbox{({\\it l,b} = 289\\fdg47, +0\\fdg12)} (clump A)  projected near the border of \\ngc, which is physically related to the nebula. Radio continuum images suggest that  the highest electron density area of the \\hii region (coincident with the ionization front)     is compressing  the densest part of the molecular overdensity,  while the low electron density region  is undergoing a champagne phase.   Three MSX compact  \\hii region (C\\hii) candidates were also reported lying at the inner border of the nebula, which confirm the mentioned scenario (see Fig. \\ref{fig:ir-cont}). In spite of the strong evidences of interaction between \\ngc\\ and its molecular environment,  disparities in angular resolution  made it difficult a direct comparison  between molecular and radio continuum/IR images.  Furthermore, the low angular resolution of the CO data did not allow us to detect any substructure in the molecular gas associated with \\ngc, and therefore a detailed analysis could not be done adequately.  The analysis of the molecular gas and dust associated with \\ngc\\ presents an important opportunity to study the interactions between \\hii regions with molecular clouds, and how the formation of massive stars (Pis 17) in the cloud may  affect further (or ongoing) star formation.   In this paper, we present new  \\cob,  \\coc, \\cod,  and   \\hcn\\  observations  around the \\hii\\ region \\ngc\\  carried out with the 12-m APEX\\footnote{APEX is a collaboration between the Max-Planck-Institut fur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory} telescope.  To account for the  properties of  the  dust  in the nebula and its surroundings, we use unpublished images  at 870 $\\mu$m from the ATLASGAL survey and available IRAC-GLIMPSE images. The aim of this work is to investigate in detail the spatial distribution and physical characteristics of the molecular gas and dust  associated with \\ngc, and to compare them with the different gas components observed with similar angular resolutions. This analysis will also allow us to identify regions of   dense molecular gas where star formation may be developing.  To look for signatures of star formation,  a new  discussion of YSO candidates in the surroundings of  \\ngc\\    using WISE and 2MASS data is also presented. In Fig. \\ref{fig:ukst}   we show the UKST red image\\footnote{http://www-wfau.roe.ac.uk/sss/index.html} of the area of the nebula NGC3503 surveyed by our molecular observations ",
        "conclusions": "Using  APEX \\cob, \\coc, \\cod, and \\hcn\\ line data, and ATLASGAL 870 $\\mu$m images, we carried out a multifrequency study of the molecular gas and dust associated with the \\hii\\ region/star forming region  \\ngc. To analyze the star formation process in the region we  made use of the WISE and 2MASS data obtained from the IPAC archive. This work is a follow-up study of \\citet{dvcca12}. The main results can be summarized as follows: \\begin{enumerate}  \\item The ionized gas of \\ngc\\     is expanding against the molecular gas component in the velocity range from $\\sim$ $-$28 to $-$23 \\kms. The morphology of the molecular gas close to  the nebula, the location of the PDR, and the  shape of radio continuum emission  confirm the ``champagne flow'' scenario proposed in Paper I.  \\item New APEX observations allowed the small scale structure of the molecular gas associated with the nebula (previously reported in Paper I) to be fully imaged. We identified several  molecular clumps (C1 to C11) and studied their physical and dynamical properties to investigate the impact of  the expanding nebula and/or the southern bright rimmed cloud \\sfo\\ onto the molecular gas.  \\item We found that warmer clumps (C3, C4, and C7) are close to the \\hii\\ region, which is indicative of an  external heating source, most probably by photoionization of their  surface molecular layers by the intense UV field of Pis 17. Warmer clumps are also denser, which suggests that they are submitted to an external compression due to the expansion of \\ngc. These clumps are likely molecular gas of the parental cloud that has been  swept up by the expansion of the ionization front and has been condensed. Clump  C4 is also adjacent to \\sfo, which might explain its highest temperature and density. A noticeable velocity gradient is  detected in C4.  This gradient may  indicative of a kinematical disturbance, or else, be the consequence of two molecular cores at slightly different velocity that are not fully resolved in the carbon monoxide emissions.  \\item Clumps located near to \\ngc\\ also have the highest LTE and virialized masses. They also exhibit the highest $M_{\\rm LTE}/M_{\\rm vir}$\\ ratio (predominantly $>$1), which,  according to the classical interpretation, indicates that they are gravitationally bound.  \\item  All molecular clumps exhibit line broadening  beyond the thermal width, which possibly indicates that  in some cases the line widths of the composite spectrum of the clumps have different velocities,  rather than showing turbulence effects. Also, spatially unresolved process like ouflows, infall, accretion, etc.  might be contributing in the broadening.  \\item We have analyzed the 870 $\\mu$m emission, characteristic of filaments and dense molecular cores.  We detected emission only in the direction of clumps C3, C4, and C7. This is also indicative of high density gas.  \\item We have presented some evidence of stellar formation in the region by detecting sources with IR excess. Unlike 2MASS candidate YSOs, WISE Class-I and Class-II candidates show a spatial distribution concentrated to the IR nebula. Several sources are detected along the external border of two of the densest molecular clumps (C4 and C7) which  suggests that they might be formed in  the compressed layers of molecular gas. Three sources are projected close to  HCN, C$^{18}$O, and 870 $\\mu$m emission peaks (coincident with the position  of clumps C4 and C3). This very likely indicates that they were born inside high density  molecular cores, making clumps C4 and C3 excellent candidates to further investigate star formation with higher spatial resolution instruments like ALMA.  \\item Since the dynamical age and fragmentation time derived for  the molecular layer differ to the age and radius of the nebula (Paper I) we have excluded the ``collect and collapse'' scenario for the YSO formation. Instead, we have proposed here some alternative mechanisms, such as ``radiative-driven implosion''  in pre-existing gravitationally bound clumps (C3, C4, and C7), or small-scale Jeans gravitational instabilities in the swept-up layers of molecular gas.    \\end{enumerate}"
    },
    "1402/1402.5088_arXiv.txt": {
        "abstract": "We report the discovery of the new pulsar wind nebula (PWN) G141.2$+$5.0 in data observed with the Dominion Radio Astrophysical Observatory's Synthesis Telescope at 1420~MHz. The new PWN has a diameter of about 3.5\\arcmin, which translates at a distance of 4.0~kpc to a spatial extent of about 4~pc. It displays a radio spectral index of $\\alpha\\approx -0.7$, similar to the PWN G76.9$+$1.1. G141.2$+$5.0 is highly polarized up to 40\\,\\% with an average of 15\\,\\% in the 1420 MHz data. It is located in the centre of a small spherical HI bubble, which is expanding at a velocity of 6~km\\,s$^{-1}$. The bubble's systemic velocity is $-53$~km\\,s$^{-1}$ and could be the result of the progenitor star's mass loss or the shell-type SNR created by the same supernova explosion in a highly advanced stage. The systemic velocity of the bubble shares the velocity of \\ion{H}{1} associated with the Cygnus spiral arm, which is seen across the 2nd and 3rd quadrants and an active star-forming arm immediately beyond the Perseus arm. A kinematical distance of 4$\\pm$0.5~kpc is found for G141.2$+$5.0, similar to the optical distance of the Cygnus arm (3.8$\\pm$1.1~kpc). % G141.2$+$5.0 represents the first radio PWN discovered in 17 years and the first SNR discovered in the Cygnus spiral arm, which is in stark contrast with the Perseus arm's overwhelming population of shell-type remnants. ",
        "introduction": "Pulsar wind nebulae (PWNe) are created by fast spinning neutron stars that produce highly relativistic magnetized particle winds, inflating an expanding bubble, confined by the expanding ejecta of the supernova explosion. Electrons and positrons are accelerated to relativistic speeds at the termination shock and interact with the magnetic field generating a synchrotron emitting nebula, called a pulsar wind nebula.  A PWN is characterized by a central concentration of radio continuum emission that gradually declines to the outer edge, in contrast to the more common shell-type supernova remnants (SNRs), which have a limb-brightened appearance. PWNe typically have flat radio continuum spectra with a spectral index $\\alpha$ between 0.0 and $-$0.3 ($S\\sim\\nu^\\alpha$, $S$: flux density, $\\nu$: frequency), although steeper spectra have been observed for PWNe DA495 \\citep{koth08b} and G76.9$+$1.1 \\citep{land93,koth06}. Many of these PWNe, in particular most of the very bright objects like the Crab nebula and 3C58, do not show an accompanying shell-type SNR. This is very puzzling since typically they are relatively young so that the actual remnant of the supernova should not have dissipated, yet.  The Canadian Galactic Plane Survey \\citep[CGPS,][]{tayl03} has portrayed the outer Galactic plane with excellent sensitivity and fidelity to extended emission in radio continuum and HI. These qualities have led to the discovery of 13 supernova remnants (SNRs) over the past 12 years \\citep{koth01b,koth03b,koth05,tian07,fost13,gerb13}. ",
        "conclusions": "G141.2+5.0 is a new radio pulsar wind nebula discovered in follow-up observations of the Canadian Galactic Plane Survey, and the first radio PWN discovered since G63.7$+$1.1 was reported by \\citet{wall97} some 17 years ago. The PWN is a resident of the Cygnus arm about 4~kpc away and is the first SNR discovered in that spiral arm. It displays a steep radio continuum spectrum and has a high level of linear polarization. It sits inside an expanding HI shell, which very likely is related to the expanding shell-type remnant of the same supernova explosion at a highly advanced phase of evolution. This may identify pulsar wind nebulae without an accompanying shell-type remnant as the result of a low energy supernova explosion."
    },
    "1402/1402.0664_arXiv.txt": {
        "abstract": "We present the results of the longest yet undertaken search for apparently extragalactic radio bursts at the Bleien Radio Observatory covering 21000 hours (898 days). The data were searched for events of less than 50 ms FWHM duration showing a $\\nu^{-2}$ drift in the spectrogram characteristic of the delay of radio waves in plasma. We have found five cases suggesting dispersion measures between 350 and 400 cm$^{-3}$ pc while searching in the range of 50 -- 2000 cm$^{-3}$ pc. Four of the five events occurred between 10.27 and 11.24 a.m. local civil time. The only exception occurred at night with the full moon in the beam. It was an event that poorly fits plasma dispersion, but had the characteristics of a solar type III burst. However, we were not able to confirm that it was a lunar reflection. All events were observed with a log-periodic dipole within 6800 hours but none with  a more directional horn antenna observing the rest of the time. These properties suggest a terrestrial origin of the \"peryton\" type reported before. However, the cause of these events remains ambiguous. ",
        "introduction": "Short-duration broadband radio pulses have been predicted from several extragalactic processes, including giant pulses from pulsars \\citep{2004ApJ...612..375C}, prompt emission of gamma-ray bursts \\citep{1993ApJ...417L..25P} and evaporating mini black holes \\citep{1977Natur.266..333R}. Such impulsive emission is delayed in a plasma relative to the speed of light. The delay is proportional to the traversed free electron column density, known as the dispersion measure (DM), and to $\\nu^{-2}$ where $\\nu$ is the observing frequency. A simultaneously emitted broadband pulse would thus show a characteristic drift in time and frequency, lower frequencies arriving later.  A number of previous searches for dispersed single radio pulses were negative. Particular galaxies were searched without detection of giant pulses \\citep[e.g.][]{2003ApJ...596.1142C}, although  \\citet{2013MNRAS.428.2857R} report possible candidates from M31 having a DM of 54.7 pc cm$^{-3}$. Other searches were aimed at detecting radio pulses associated with gamma-ray bursts. No radio event was found by \\citet{1998A&A...329...61B}, but \\citet{2012ApJ...757...38B} report two possible cases delayed from the gamma-ray peak by several minutes. Several general searches not aimed at particular objects were reported without detections. Amy et al. (1989) operated a transient event monitoring system for a total of 180 days at 843 MHz in parallel to interferometric observations at the Molonglo Observatory. \\citet{2003PASP..115..675K} did not observe any non-solar event within 18 months with the STARE system of three geographically separated instruments at 611 MHz. \\citet{2012ApJ...744..109S} did not detect any extragalactic pulses within 450 hours of a search with the Allen Telescope at 1430 MHz.  Recently four promising detections were reported from a general search with a multi-beam receiver of the Parkes telescope at 1358 MHz \\citep{2013Sci...341...53T}. The events were detected in only one beam of 13, suggesting a non-terrestrial origin. The range of dispersion measures was from 553 to 1103 pc cm$^{-3}$, the FWHM duration from $<$1.1 to 5.6 ms. The total observing time of the 13 beams amounted to 23 days each. In the various Parks searches some 25 other pulses drifting as $\\nu^{-2}$ were detected in all beams of the receiver, thus seem to have come into the telescope through sidelobes (Burke-Spolaor et al. 2011). They are known as \"perytons\". These events show apparent DMs of 350 - 400 pc cm$^{-3}$, but some deviate from the strict dispersion delay. The reported fluxes are up to 272 kJy, and the FWHM durations are some 20 ms or more. An event reported earlier by Lorimer et al. (2007) was observed in several beams, lasted only 4.6 ms, but may be of a similar kind. Five more events just before a group of Burke-Spolaor events were reported by \\citet{2012MNRAS.420..271K} and another four by \\citet{2012MNRAS.425.2501B}. Out of the 25 perytons reported so far, a group of 16 occurred within 30 minutes. Another two were found within one minute. Thus the number of independent events is 9, all observed at Parkes Observatory in New South Wales, Australia. Atmospheric conditions such as heavy rain or terrestrial X-ray and gamma-ray events were found not to be associated \\citep{2012MNRAS.425.2501B}. If man-made, these signals traverse the globally protected frequency band at 1420 MHz (21 cm). Perytons are generally assumed to be terrestrial, but their origin is unknown and they can easily be confused with extragalactic pulses.  Here we report on long-term observations with small antennas having a large field of view but a short range in distance. ",
        "conclusions": "\\label{sect:cl}  Five events of plasma-dispersion like drifts in frequency have been observed in an observation of several years. These events have the following properties:  \\begin{itemize} \\item The detection rate is much higher than the one-beam events reported by \\citet{2013Sci...341...53T}. \\item The duration is of the order of 20 ms. \\item The dispersion measure is in the range 350 - 400 cm$^{-3}$pc. \\item The bursts vary strongly in intensity at different frequencies. \\item Four out of five bursts occurred between 10.27 and 11.24 a.m. local civil time. \\item One out of five events deviates from the strict dispersion delay. \\item All events were observed with the low-directional log-periodic antenna. \\item No solar events were observed at the time of the events. \\item No laboratory equipment was operated at the time of the events. \\item No lightning or other atmospheric phenomena were recorded on site at the time of the events. \\end{itemize}  The burst characteristics (first 6 points) listed above are consistent with the many-beam events observed with the Parks Telescope, thus showing presumably terrestrial origin \\citep{2011ApJ...727...18B}. These authors created the name \"perytons\" to describe their deceiving similarity with extragalactic events.  Our events were independent and well separated in time. Their rate is about one per 57 days observed with the log-periodic antenna. These events are the first ones of this kind observed outside of the Parks Observatory region. Their origin remains unexplained."
    },
    "1402/1402.3983_arXiv.txt": {
        "abstract": "{} {Our  aim is  to show  how variable  Doppler boosting of an intrinsically variable  jet can explain the long-term  modulation of $1667 \\pm 8$\\,days observed in the radio emission of  \\lsp.} {The  physical scenario is that of a conical, magnetized plasma jet having a periodical ($P_1$) increase of relativistic particles, $N_{\\rm rel}$, at a specific orbital phase, as predicted by  accretion  in the eccentric orbit of \\lsp. Jet precession ($P_2$) changes the angle, $\\eta$, between jet axis and line of sight, thereby inducing variable Doppler boosting. The problem is defined in spherical geometry, and  the optical depth through the precessing jet is calculated by taking  into account that the plasma is stratified along the jet axis. The synchrotron emission of such a jet was calculated  and we fitted the resulting flux density $S_{\\rm model}(t)$   to the observed flux density  obtained during a 6.5-year  monitoring  of \\lsi by the  Green Bank radio interferometer.} {Our physical model for the system \\lsi  is not only able to reproduce the long-term  modulation in the radio emission, but it  also reproduces all the other observed characteristics of the radio source, the orbital modulation of the outbursts, their orbital phase shift, and their  spectral index properties. Moreover, a correspondence seems to exist between variations in the ejection angle  induced by precession and the  rapid rotation in position angle observed in  VLBA images. } {The peak of the long-term modulation occurs when the jet electron density is around its maximum  and the approaching jet is forming  the smallest possible angle with the line of sight. This coincidence of maximum number of emitting particles and  maximum Doppler boosting of their emission occurs every $\\sim$1667 days and creates the long-term modulation observed in \\lsp.} ",
        "introduction": "A recent timing analysis  of the GBI radio data of  \\lsi has revealed  two  frequencies: $P_1={1\\over \\nu_1}=\\unit[26.49 \\pm 0.07]{days}$ and $P_2={1\\over \\nu_2}=\\unit[26.92 \\pm 0.07]{days}$ \\citep{massijaron13}. The  period $P_1$ agrees with the value of $\\unit[26.4960\\pm0.0028]{days}$ \\citep{gregory02} associated to the  orbital period  of the binary system formed by a compact object and a Be star \\citep{grundstrom07}. Period  $P_2$  agrees with the previous estimate of 27-28 days  by radio astrometry of the precessional period of the radio structure \\citep{massi12}.  The timing analysis results seem to confirm   \\lsi  as the  second case of a radio-emitting X-ray binary with     precessional period ($P_2$) of its  radio jet very close to the  orbital ($P_1$) period of the binary system \\citep{massi12, massijaron13}. Whereas the several  known precessing X-ray binaries \\citep{larwood98} have a precession period an order of magnitude longer than the orbital period (as predicted for  a tidally forced precession by the companion star),  the different case of   GRO J1655-40 was discovered in   1995  \\citep{hjellmingrupen95}. This black hole candidate, which has an orbital period of   2.601$\\pm 0.027$~days \\citep{bailyn95}, revealed a radio jet with a  precessional period  of 3.0$\\pm$0.2 days  \\citep{hjellmingrupen95}. \\lsi is a high-mass X-ray binary, and  \\gro{} is a low-mass X-ray binary. The mechanism proposed in both cases to explain the precession is based on general relativity effects around the compact object. Lense-Thirring precession has been analysed for \\gro{} by \\citet{martin08} and for \\lsi{} by \\citet{massizimmermann10}.  When  two frequencies are only slightly different, a beating results,  i.e. an interference producing a  new frequency:    their average, $ \\nu_{\\rm average}={{\\nu_1+\\nu_2}\\over 2}$, which gets modulated with  the beat frequency, $\\nu_{\\rm beat}=\\nu_1-\\nu_2$. As shown  in \\citet{massijaron13} for \\lsi, the term $1/\\nu_{\\rm beat}$ results  $\\unit[1667 \\pm 393]{days}$ and is compatible with  the observed long-term modulation of the radio flux density of $1667 \\pm 8$\\,days, attributed in the past to variations in the wind of the Be star \\citep{gregoryneish02}. Evidence of  beating in \\lsi   comes from the fact that the periodicity of the radio outbursts is indeed the one predicted  by the beat theory, which is $P_{\\rm outburst}=P_{\\rm average}={1\\over\\nu_{\\rm average}}=26.70\\pm 0.05$ days \\citep{massijaron13, jaronmassi13}. Since the the two frequencies $ \\nu_{\\rm average}$ and  $\\nu_{\\rm beat}$ are related , the long-term modulation should also be a result of the beating process between $P_1$ an $P_2$.  The aim of  this paper is  to prove, in the context of a physical scenario, how the mutual relationship between $P_1$ and $P_2$ can generate the long-term periodicity. We linked the two periodicities, $P_1$ and $P_2$, to two physical processes: the periodical ($P_1$) increase  of relativistic electrons in a conical  jet and  the periodical ($P_2$)  Doppler boosting of the emitted radiation by relativistic electrons because of jet precession. The hypothesis that  the compact object in \\lsp, accreting  material from the  Be wind, undergoes a  periodical  ($P_1$) increase in the accretion  at a  particular orbital phase along an eccentric orbit was suggested and developed  by several authors  \\citep{taylor92, martiparedes95, boschramon06, romero07}. The hypothesis that a precessing jet, with an approaching jet having  large excursions in its position angle, should give rise to appreciable Doppler boosting effects is supported by the morphology of \\citet{massi04}, \\citet{dhawan06}, and \\citet{massi12} images  showing extended radio structures changing from two-sided to one-sided morphologies at  different position angles.  Our model describes a precessing ($P_2$) jet with a periodically ($P_1$) varying number of relativistic electrons. Section~2 describes our analytical model and how to calculate its flux density $S_{\\rm model}$. In Sect. 3, all parameters  for calculating  $S_{\\rm model}$ are derived by fitting the model to  6.5 years of GBI radio flux density observations. In Sect.\\,4 we present our results and in Sect.\\,5 our conclusions.  \\begin{figure}[ht] \\centering \\includegraphics[scale=.3, clip]{Fig1.pdf}\\\\ \\caption{Precessing jet and relative angles. Left: Sketch of the precessing cone. The angle $\\zeta$ is the angle between the jet precession axis and the line of sight (l.o.s.)  direction, $\\psi$ is the opening  angle of the precession cone and $\\eta$ is the continuosly changing angle between  the jet and the line of sight because of precession. Right: The  angles $\\zeta$,$\\eta$,$\\psi$ in spherical geometry. The angle $\\Omega$, describing in the precession period $P_2$ an angle  $2\\pi$,  defines the temporal dependence. } \\label{Fig1} \\end{figure} ",
        "conclusions": "In this paper we have developed a model that reproduces the observed long-term modulation of the   radio emission of the binary system \\lsi. The hypothesis is that the radio emission is due to a jet precessing with a period of $P_2=\\unit[26.92 \\pm 0.07]{days}$. The jet is  anchored on the accretion disk of a compact object orbiting, with a period of $P_1=\\unit[26.49 \\pm 0.07]{days}$, the Be star. Variable accretion around the eccentric orbit increases the  relativistic electron density in the jet at a particular orbital phase. We fitted  our theoretically derived  flux density, $S_{\\rm model}(t)$  (see Eq. 1)  to 6.5 yr GBI radio observations and determined  the parameters of Table 1.  We found that the  precessing ($P_2$) jet model with intrinsic periodical variations ($P_1$) of emitting particles reproduces all seven  observational facts: \\begin{enumerate} \\item the  long-term modulation  \\citep[and references there]{gregory02}, see Fig. \\ref{mod}\\,a and Fig. \\ref{shift}\\,bottom; \\item the  orbital shift of the peak of the radio outburst \\citep{paredes90}, see Fig. \\ref{shift}\\,bottom; \\item the periodical outburst with orbital occurrence \\citep[and references there]{gregory02}, see Figs. \\ref{mod}\\,b and c and  Fig. \\ref{beat}\\,a; \\item the   double clustering for data folded with $P_1+P_2\\over 2$  \\citep{massijaron13}, see Fig. \\ref{beat}\\,c; \\item the  spectral index evolution \\citep{massikaufman09}, see Fig. \\ref{index}; \\item a variation of the ejection angle suggestive of the  rapid rotation in position angle of VLBA maps \\citep{dhawan06,massi12}, see Fig. \\ref{vlba}; \\item the strength of the magnetic field \\citep{massi93}, see Sect. 4.6. \\end{enumerate}  The maximum of the long-term modulation occurs when the maximum ejection of relativistic electrons takes place  and  at  the same time the  approaching jet is forming  the smallest possible angle with the line of sight (i.e., the Doppler boosting is at its maximum). After a cycle of duration $P_2$, when the approaching jet again forms the smallest angle $\\eta$, the  ejection is now already on its decay phase (being $P_1 < P_2$) and its Doppler boosting gives rise to an outburst with a lower amplitude than that of one cycle before. As we quantitatively show in Sect. 4.2, the peak of the outburst gets lower and lower at each cycle until the most unfavorable situation occurs, i.e., when the jet is pointing towards us while the main ejection is at its minimum. After this point, it happens that when the approaching jet again forms the smallest $\\eta$ angle, there is the beginning of a new ejection of particles and the observed flux density slowly rises again. The period of 1667$\\pm$8 days \\citep{gregory02} of the long-term modulation is then just the number of $P_1$ cycles needed to compensate for the difference $P_2-P_1$ and synchronize $P_1$  again to  $P_2$, which is just $[P_2/(P_2-P_1)]P_1$ \\citep{massijaron13}.  Our  model,   meant in the first instance  to explain the long-term modulation affecting the radio emission of \\lsp, can become a powerful tool for understanding the physical processes at work in this gamma-ray binary. For this purpose the model should be   extended to include internal shocks in the conical flow, i.e., the transient jet. As  discussed in Sect. 4.4. the transient jet observed in \\lsi is associated to a second  population of relativistic electrons, with $p=2.0$, which is therefore different from the population with $p=1.8$,  forming the conical jet analysed in this paper. It is quite interesting that two similar  populations are invoked to explain emission in the  TeV  and GeV band  \\citep{zabalza11}. The important implications, if confirmed, are that the GeV emission is related to inverse Compton of the electrons of the conical jet and that the TeV emission is related to  electrons of the transient jet.  This model, with conical jet and transient jet can be tested by fitting it  to the radio spectral  index data analysed in \\citet{massikaufman09}. The  theory of accretion in  an eccentric orbit  (Sect. 2.4) also predicts an ejection around periastron passage, along with the ejection at orbital phase $\\Phi$=0.6 considered here  (i.e., towards apoastron). By fitting the model to the  radio spectral  index data, one could  verify if  each of the two populations of electrons is  generated twice along the orbit, i.e., that the transient jet occurs at periastron as well.  Indeed, at  different epochs, emission towards  apoastron \\citep{albert09} and  periastron \\citep{acciari11} was detected by Cherenkov telescopes. Different  Doppler boosting of the two ejections along the eccentric orbit could   explain the observed variations. Maximum Doppler boosting occurring between the two ejections could explain the broad gamma-ray emission without any clear orbital modulation observed in Fermi-LAT data \\citep{hadasch12, ackermann13}.  Finally, considering that one contribution to the gamma-ray emission consists of upscattered stellar photons, the role of Doppler boosting in  gamma-ray emission is probably  even  more important than in the radio band. In fact,  the amplification due to the Doppler factor for external inverse Compton is higher than for synchrotron emission \\citep[and reference therein]{kaufman02}."
    },
    "1402/1402.6618_arXiv.txt": {
        "abstract": "In recent years, there have been several successful attempts to constrain the equation of state of neutron star matter using input from low-energy nuclear physics and observational data. We demonstrate that significant further restrictions can be placed by additionally requiring the pressure to approach that of deconfined quark matter at high densities. Remarkably, the new constraints turn out to be highly insensitive to the amount --- or even presence --- of quark matter\tinside the stars. ",
        "introduction": "\\label{sec:intro}  The equation of state (EoS) of cold and dense strongly interacting matter, which determines the inner structure of compact stars \\citep{Glendenning:1997wn}, is encoded in its fundamental theory, Quantum Chromodynamics (QCD). A full nonperturbative determination of the pressure of the theory is still out of reach due to the so-called Sign Problem of lattice QCD \\citep{deForcrand:2010ys}. Nevertheless, the methods of chiral effective field theory (EFT) of nuclear forces \\citep{Epelbaum:2008ga} and high-density perturbative QCD (pQCD) \\citep{Kraemmer:2003gd} have matured enough to provide reliable predictions for the EoS in the limits of low density nuclear matter and dense quark matter, respectively. In particular, by now both approaches produce results with reliable error estimates, implying that it is finally possible to quantitatively estimate our understanding of the neutron star matter EoS.  During the past couple of years, several articles have addressed the determination of the neutron star EoS by combining insights from low-energy chiral EFT with the requirement that the resulting EoSs support the most massive stars observed (see e.g.~\\cite{Hebeler:2013nza}). In particular, the discovery of neutron stars with masses around two solar masses \\citep{Demorest:2010bx,Antoniadis:2013pzd} has recently been seen to lead to strong constraints on the properties of stellar matter \\citep{Lattimer:2012nd}. While otherwise impressive, these analyses have solely concentrated on the low density regime, and have typically applied no microphysical constraints beyond the nuclear saturation density $n_0$. This has resulted in EoSs that behave very differently from that of deconfined quark matter even at rather high energy densities.  In the present work, our aim is to demonstrate that the EoS of neutron star matter can be significantly further constrained by requiring it to approach the quark matter one at high density. To do this, we use the state-of-the-art result of \\cite{Fraga:2013qra}, where a compact expression for the three-loop pressure of unpaired quark matter, taking into account the nonzero value of the strange quark mass, was derived (see also \\cite{Kurkela:2009gj} and \\cite{Kurkela:2010yk} for details of the original pQCD calculation). A particularly powerful outcome of the analysis is that the high density constraint significantly reduces the uncertainty band of the stellar matter EoS even at low densities, well below a possible phase transition to deconfined quark matter. This implies that the $M$-$R$ relations we obtain are more restrictive than previous ones even for pure neutron stars.  In practice, our calculation proceeds as follows (see also fig.~\\ref{fig1}): at densities below $1.1n_0$, we employ the chiral EFT EoS of \\cite{Tews:2012fj}, assuming the true result to lie within the error band given in this reference. At baryon chemical potentials above 2.6 GeV, where the relative uncertainty of the quark matter EoS is as large as the nuclear matter one at $n=1.1 n_0$, we on the other hand use the result of \\cite{Fraga:2013qra} and its respective error estimate. Between these two regions, we assume that the EoS is well approximated by an interpolating polytrope built from two ``monotropes'' of the form $P(n)$=$ \\kappa n^\\Gamma$. These functions are first matched together in a smooth way, but later we also consider the scenario of a first-order phase transition, allowing the density to jump at the matching point of the two monotropes.  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig1.eps} \\caption{Known limits of the stellar EoS on a logarithmic scale. On the horizontal axis we have the quark chemical potential (with an offset so that the variable acquires the value $0$ for pressureless nuclear matter), and on the vertical axis the pressure. The band in the region around the question mark corresponds to the interpolating polytropic EoS that will be introduced in this work.} \\label{fig1} \\end{figure}  Varying the polytropic parameters and the transition density over ranges limited only by causality, we obtain a band of EoSs that can be further constrained by the requirement that the EoS support a two solar mass star. This results in EoS and $M$-$R$ bands that are significantly narrower than ones obtained without the high density constraint (see Fig.~\\ref{fig10}). An important check of the robustness of our construction is that the obtained band is largely unaffected by the nature of the assumed phase transition or by the introduction of a third interpolating monotrope.   \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig2.eps} \\caption{The interpolated pressure of nuclear and quark matter, normalized by the pressure of a gas of free quarks and shown together with the pQCD result at high densities. All generated EoSs lie within the shaded green and turquoise areas, of which only the green ones support a star of $M=2M_{\\odot}$. Three representative EoSs marked with I-III have crosses denoting the maximal chemical potential reached at the center of the star.} \\label{fig2} \\end{figure}  Clearly, both our setup and results bear some resemblance to those of \\cite{Hebeler:2013nza}. An important difference between these two calculations is, however, that while the authors of the latter paper imposed a set of somewhat \\textit{ad hoc} constraints on their polytropic parameters, for us this is not necessary, as the high density constraint automatically restricts these numbers. In this vein, one can in fact argue that our calculation gives an \\textit{a posteriori} justification for many of the choices made in \\cite{Hebeler:2013nza}. At the same time, it is also important to note that the ease with which we are able to perform the matching between the low and high density EoSs --- and the high maximum masses we obtain for the stars --- is in stark contrast with many earlier attempts to directly match nuclear and quark matter EoSs onto each other (cf.~e.g.~\\cite{Fraga:2001xc} and \\cite{Alford:2004pf}).  Our paper is organized as follows. In section 2, we first explain the details of our calculation, i.e.~introduce the low and high density EoSs used as well as the specific parametrization of our interpolating polytropes. After this, we proceed to display and analyze our results in section 3, covering both the EoSs, $M$-$R$ relations, and various correlations between the parameters appearing in our calculation. In section 4, we finally draw our conclusions. ",
        "conclusions": "In the paper at hand, we have constructed a novel scheme for determining the EoS of compact star matter that involves an interpolation between the regimes of low-energy chiral effective theory and high-density perturbative QCD. These two limiting results are truly robust within their ranges of applicability, as they represent controlled calculations in the fundamental theory of the strong interactions. Our work on the other hand constitutes the first ever attempt to take constraints from both sides on equal footing when determining the EoS between these limits. We have demonstrated that this leads to important new constraints on the properties of compact star matter on a wide density range, and thus even for stars containing only hadronic matter.  The strictness of the constraints placed on the stellar EoS by its high-density limit can be understood through the tension between the softness of the perturbative EoS and the stiffness required by the confirmed existence of a two solar mass compact star. For the two interpolating monotropes we employ in our calculation, this translates into a significant difference between the respective polytropic indices: While the first one needs to be rather stiff, with $\\gamma_1>2.86$, the latter must be considerably softer, $1<\\gamma_2<1.5$. Although the polytropes themselves of course do not carry information about the underlying microphysics, such a strong shift in the polytropic index might be interpreted as a sign of the effective degrees of freedom of the system changing from hadronic to denconfined ones.  The effect of the high density constraint is perhaps best illustrated in Fig.~\\ref{fig10}, which displays our EoS band in the form of energy density vs.~pressure, plotted together with the previous prediction of \\cite{Hebeler:2013nza}, dubbed HLPS. The latter work applied the same low-density EoS we did and took into account the two solar mass constraint, but did not require the result to approach the pQCD EoS at large densities. As expected, the main difference between the two results is seen in the HLPS cloud containing somewhat softer EoSs at low density and stiffer ones at high density.  The rather narrow EoS band that results from our interpolation naturally corresponds to a well defined region in the mass-radius diagram of compact stars. For a $1.4M_\\odot$ neutron star, the radii we obtain range between 11 and 14.5$\\,$km, while the radius of a $2M_\\odot$ pulsar lies within $R\\approx 10-15\\,$km. Interestingly, we do not find configurations with masses above $2.75M_\\odot$ (for bitropic interpolation the maximal mass is $2.5M_\\odot$). This conclusion is in contrast with what has been found before without the high-density constraint; see e.g.~\\cite{Hebeler:2013nza}, where stars with masses up to $3M_\\odot$ were discovered.  For the convenience of the reader, we finally provide three representative EoSs in a tabulated form at the end of this paper. These EoSs are all subluminal, able to sustain a two solar mass star, and maximally different from each other. Of them, EoS I gives the minimal radius, EoS II the maximal mass and EoS III the maximal radius for our compact stars.  In conclusion, we find it remarkable, how the properties of quark matter at asymptotically high densities can be seen to have such a strong impact on the structure of compact stars at much lower energies. As we have highlighted in Fig.~\\ref{fig1}, this fact appears to make it possible to largely bridge the gap between the respective EoSs of low-density nuclear matter and high-density (perturbative) quark matter."
    },
    "1402/1402.3517_arXiv.txt": {
        "abstract": "In the quest to understand the formation of the building blocks of life, amorphous solid water (ASW) is one of the most widely studied molecular systems. Indeed, ASW is ubiquitous in the cold interstellar medium (ISM), where ASW-coated dust grains provide a catalytic surface for solid phase chemistry, and is believed to be present in the Earth's atmosphere at high altitudes. It has been shown that the ice surface adsorbs small molecules such as CO, N$_2$, or CH$_4$, most likely at OH groups dangling from the surface. Our study presents completely new insights concerning the behaviour of ASW upon selective infrared (IR) irradiation of its dangling modes. When irradiated, these surface H$_2$O molecules reorganise, predominantly forming a stabilised monomer-like water mode on the ice surface. We show that we systematically provoke ``hole-burning'' effects (or net loss of oscillators) at the wavelength of irradiation and reproduce the same absorbed water monomer on the ASW surface. Our study suggests that all dangling modes share one common channel of vibrational relaxation; the ice remains amorphous but with a reduced range of binding sites, and thus an altered catalytic capacity.\\\\ ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.6110_arXiv.txt": {
        "abstract": "{} {We want to estimate whether the positrons produced by the $\\beta^{+}$-decay of \\alu,\\titan\\ and \\nickel\\ synthesised in massive stars and supernovae are sufficient to explain the 511 \\kev\\ annihilation emission observed in our Galaxy. Such a possibility has often been put forward in the past. In a previous study, we showed that nucleosynthesis positrons cannot explain the full annihilation emission. Here, we extend this work using an improved propagation model.} {We developed a Monte Carlo Galactic propagation code for $\\sim$MeV positrons in which the Galactic interstellar medium, the Galactic magnetic field and the propagation are finely described. This code allows us to simulate the spatial distribution of the 511 \\kev\\ annihilation emission. We test several Galactic magnetic fields models and several positron escape fractions from type-Ia supernova for \\nickel\\ positrons to account for the large uncertainties in these two parameters. We consider the collisional/ballistic transport mode and then compare the simulated 511\\kev\\ intensity spatial distributions to the INTEGRAL/SPI data.} {Whatever the Galactic magnetic field configuration and the escape fraction chosen for \\nickel\\ positrons, the 511 \\kev\\ intensity distributions are very similar. The main reason is that $\\sim$MeV positrons do not propagate very far away from their birth sites in our model. The direct comparison to the data does not allow us to constrain the Galactic magnetic field configuration and the escape fraction for \\nickel\\ positrons. In any case, nucleosynthesis positrons produced in steady state cannot explain the full annihilation emission. The comparison to the data shows that: (a)  the annihilation emission from the Galactic disk can be accounted for; (b) the strongly peaked annihilation emission from the inner Galactic bulge can be explained by positrons annihilating in the central molecular zone but this seems to require more positron sources than the population of massive stars and type Ia supernovae usually assumed for this region;  (c) the more extended emission from the Galactic bulge cannot be explained. We show that a delayed 511 \\kev\\ emission from a transient source, such as a starburst episode or a recent activity of Sgr A*, occurring between 0.3 and 10 Myr ago and producing between 10$^{57}$ and 10$^{60}$ sub-MeV positrons could explain this extended component, and potentially contribute to the inner bulge signal. } {} ",
        "introduction": "\\label{intro}  The 511 keV line emission from our Galaxy is unambiguously produced by low energy positrons that annihilate with electrons, but the exact origin of these positrons remains unclear. The spatial distribution of the annihilation emission was measured by several generations of gamma-ray instruments and most recently by INTEGRAL/SPI\\footnote{For more information on SPI, see \\citet{Vedrenne03}.} \\citep{Knodlseder05,Weidenspointner08b,Bouchet10}. It comprises faint emission from the inner part of the Galactic disk (GD), and a strong diffuse emission from the Galactic bulge \\citep[GB, which can be modelled with a narrow and a wide spheroidal gaussian distribution with projected FWHM of $\\sim$3$\\degree$ and $\\sim$11$\\degree$ respectively, see][]{Weidenspointner08b}. This emission is very particular, with an inferred bulge-to-disk luminosity ratio ranging from 2 to 6. None of the known Galactic astrophysical object or interstellar matter distributions resembles the annihilation emission distribution.  In order to explain this bulge-to-disk luminosity ratio, several authors suggested that positrons produced by supernovae in the disk could propagate far enough to annihilate in the bulge \\citep{Prantzos06,Higdon09}. Other authors proposed that mini-starbursts or the supermassive black hole Sgr A* produced a large amount of positrons 10$^6$ or 10$^7$ years ago, which filled the GB and are annihilating now \\citep{Parizot05,Totani06,Cheng06,Cheng07}.  A major issue in such studies is that the propagation of positrons in the interstellar medium (ISM) is not well understood \\citep{Jean06}. In a previous detailed analysis, \\citet{Higdon09} used an inhomogeneous diffusion model, including collisional and collisionless transport of positrons, in a finely-structured ISM and found that nucleosynthesis positrons could account for all the observables of the Galactic annihilation emission. However, this model raised some criticism in the community \\mbox{\\citep[e.g.,][]{Prantzos11,Martin12}.} In the present paper, we would like to propose a different approach based on the theoretical investigation of \\citet{Jean09} on propagation mechanisms for positrons in the ISM. These authors showed that at low energy (E $\\lesssim$ 10 MeV), positrons do not interact with magnetohydrodynamic waves and propagate in the collisional mode by undergoing gyro-motion around magnetic field lines and collisions with gas particles. In such conditions, a correct treatment of positron propagation requires using representative models of Galactic magnetic field (GMF) and Galactic gas distributions. In a previous study, we performed a simulation of the propagation of positrons emitted in the decay of radioactive nuclei produced by massive stars and supernovae (\\alu, \\titan\\ and \\nickel) using a modified version of the GALPROP cosmic-ray propagation code \\citep{Martin12}. This code treats the transport of positrons as a diffusive process and uses 2D analytical distributions of the large-scale average gas density. In this framework, we showed that it is hard to explain the morphology of the annihilation emission from radioactivity positrons, which led us to the conclusions that either an additional component is needed to explain the bulge emission, or a finer modelling is required. We explored the latter option in the present work.  In this paper, we investigate the fate of positrons produced by stellar nucleosynthesis in our Galaxy with a Monte Carlo code that takes into account the transport of positrons in the collisional mode. This code simulates the injection, propagation, energy loss and annihilation of positrons taking into account spatial distributions for sources, interstellar gas and magnetic field. The results of the simulations allow us to derive sky maps, light curves and spectra of the annihilation emission as functions of the sources of positrons. Section \\ref{model} describes the Monte Carlo method and the various model components. In Section \\ref{simulations}, we present and discuss the results of the simulations. The simulated sky maps of the annihilation emission at 511 keV are compared to the INTEGRAL/SPI data. We show that type Ia supernovae (SNe Ia) cannot be the main source of positrons and that additional sources are needed to explain the measured disk and bulge emissions. The latter can be explained by a brief injection of a large amount of positrons in the central region of our Galaxy (e.g. in the central molecular zone) that occurred several Myr ago. In Section \\ref{conclusion}, we summarize our study and conclusions. ",
        "conclusions": "\\label{conclusion}  The aim of this work was to determine if nucleosynthesis positrons could explain the morphology of the 511 \\kev\\ emission observed in our Galaxy. Using a Monte Carlo code, we simulated their inhomogeneous collisional propagation and energy losses in a finely-structured ISM taking into account the Galactic magnetic field structure, the different ISM gaseous phases, particular features of the Galaxy such as the central molecular zone or the holed tilted disk, and testing two extreme ISM models.  We studied the contributions to the annihilation emission of positrons produced in massive stars (\\alu\\ and \\titan) and in SNe Ia (\\nickel). These sources have often been cited in the past as the most likely major contributors to Galactic positrons. Due to large uncertainties in the positron escape fraction from the SNe Ia ejecta and the structure of the magnetic field in the Galactic halo, we tested several escape fractions and halo magnetic field configurations to see the impact of these parameters on the annihilation emission morphology.  The different combinations of these parameters experimented in the simulations result in quite similar 511 \\kev\\ emission morphologies. The main reason is that nucleosynthesis positrons do not propagate far away from their birth sites. In any case, a very low fraction of positrons $\\lesssim$7$\\%$ manage to escape the Galaxy. The rest of the positrons only travel on average $\\sim$1 kpc and the 511 \\kev\\ intensity spatial distributions are thus strongly correlated with the source spatial distributions. Therefore, the steady state annihilation of nucleosynthesis positrons cannot account for the total annihilation emission observed in our Galaxy.  Comparison of our simulated sky maps to 8 years of INTEGRAL/SPI data confirms that nucleosynthesis positrons can explain the annihilation emission from the Galactic disk, but cannot fully account for that from the Galactic bulge. The morphology of the strongly peaked inner bulge emission could be explained by massive-stars positrons produced in the central molecular zone around the Galactic centre, with a possible contribution from positrons channelled there by a dipole halo field. However, depending on the magnetic field model, matching the observed intensity requires at least a contribution from prompt SNe Ia in the central molecular zone or more massive stars than currently, and at most the contribution from another unknown source. In any case, the emission from the outer bulge cannot be reproduced by steady state annihilation emission from nucleosynthesis positrons. We showed that a single and brief injection in the central molecular zone of a large amount of positrons, such as a starburst that occurred several Myr ago, could also explain the annihilation emission from the outer bulge. We found from our simulations that such an event occurring between 0.3 and 10 Myr ago and producing between 10$^{57}$ and 10$^{60}$ sub-MeV positrons could quantitatively explain the current emission from the outer bulge (and could also contribute to the inner bulge emission at some level). Nevertheless, detailed studies of these scenarios have to be undertaken before considering them as serious candidates to explain the complete 511 \\kev\\ annihilation emission from the Galactic bulge."
    },
    "1402/1402.3029_arXiv.txt": {
        "abstract": "The photometric characterization of M33 star clusters is far from complete. In this paper, we present homogeneous $UBVRI$ photometry of 708 star clusters and cluster candidates in M33 based on archival images from the Local Group Galaxies Survey, which covers 0.8 deg$^2$ along the galaxy's major axis. Our photometry includes 387, 563, 616, 580, and 478 objects in the $UBVRI$ bands, respectively, of which 276, 405, 430, 457, and 363 do not have previously published $UBVRI$ photometry. Our photometry is consistent with previous measurements (where available) in all filters. We adopted Sloan Digital Sky Survey $ugriz$ photometry for complementary purposes, as well as Two Micron All-Sky Survey near-infrared $JHK$ photometry where available. We fitted the spectral-energy distributions of 671 star clusters and candidates to derive their ages, metallicities, and masses based on the updated {\\sc parsec} simple stellar populations synthesis models. The results of our $\\chi^2$ minimization routines show that only 205 of the 671 clusters (31\\%) are older than 2 Gyr, which represents a much smaller fraction of the cluster population than that in M31 (56\\%), suggesting that M33 is dominated by young star clusters ($<1$ Gyr). We investigate the mass distributions of the star clusters---both open and globular clusters---in M33, M31, the Milky Way, and the Large Magellanic Cloud. Their mean values are $\\log(M_{\\rm cl}/M_{\\odot})=4.25$, 5.43, 2.72, and 4.18, respectively. The fraction of open to globular clusters is highest in the Milky Way and lowest in M31. Our comparisons of the cluster ages, masses, and metallicities show that our results are basically in agreement with previous studies (where objects in common are available); differences can be traced back to differences in the models adopted, the fitting methods used, and stochastic sampling effects. ",
        "introduction": "\\label{s:intro}  Since star clusters represent an important component of the galaxies they are associated with, studies of star clusters' stellar populations and age distributions can provide clues to the formation and evolution of their host galaxies. In addition, since populous star clusters are much more luminous than individual stars, they are usually much easier to observe and study.  At a distance of $847\\pm60$ kpc---equivalent to a distance modulus of $(m-M)_0=24.64\\pm0.15$ mag \\citep{gal04}---M33 (also known as the Triangulum Galaxy) is the third largest spiral galaxy in the Local Group of galaxies. Since the galaxy is seen relatively face-on, under an inclination of $i=56^{\\circ}\\pm1^{\\circ}$ \\citep{zeh}, it is eminently suitable for studies of its star cluster system. At present, the most comprehensive and widely used star cluster catalog is that of \\citet{sm07}, which combines data on almost all M33 star clusters published in the literature, including information on their photometry, ages, metallicities, and masses. The latest version\\footnote{http://www.mancone.net/m33$\\_$catalog/, updated in December 2010.} (henceforth SM10) includes 595 star clusters and candidates. \\citet{pl07} found 104 star clusters in {\\sl Hubble Space Telescope} ({\\sl HST})/Wide Field and Planetary Camera-2 (WFPC2) archival images, including 32 new objects based on new {\\sl HST} observations. Although their observations improved the spatial coverage of the M33 disk, this catalog is still incomplete for the entire disk. These authors found two different star cluster populations on the basis of their sample's color--magnitude diagram (CMD), including a large number of blue clusters and a smaller number of red objects. They also suggested that relatively more red clusters are found in the galaxy's outer regions.  Subsequently, \\citet{zkh} published a list of 4780 extended sources, including 3554 new cluster candidates observed with the MegaCam instrument on the 3.6 m Canada--France--Hawai'i Telescope (CFHT). However, $\\sim$60\\% of these clusters are not considered genuine owing to possible misidentifications \\citep{san09,san10}. Based on {\\sl HST}/Advanced Camera for Surveys (ACS)--Wide Field Channel (WFC) observations, \\citet{san09} presented photometry of 161 M33 star clusters, of which 115 were newly identified. Based on their CMDs, they suggested that these clusters' ages were between 0.01 and 1 Gyr, whereas their masses range from $5\\times10^3 M_\\odot$ to $5\\times10^4 M_{\\odot}$. However these authors also point out that, since their photometry is generally not sufficiently deep to detect the main-sequence turnoff (MSTO), very few of their sample clusters are older than 1 Gyr. Using MegaCam on the CFHT, \\citet{san10} identified 2990 extended sources in M33, 599 of which were new cluster candidates and 204 were previously known clusters. Based on CMD analysis, these authors suggested that the majority of the clusters have young to intermediate ages, although their sample also includes some old objects. They suggested that a possible M31--M33 interaction some 3.4 Gyr ago may have triggered an epoch of star (cluster) formation in M33.  Comparison of observational spectral-energy distributions (SEDs) with theoretical stellar population synthesis models by application of $\\chi^2$ minimization is a widely used technique to estimate ages, metallicities, reddening values, and masses of extragalactic star clusters. This technique has been applied to the cluster systems in, e.g., M31 \\citep{jiang03,fan06,fan10,ma07,ma09,wang10}, M33 \\citep{ma01,ma02a,ma02b,ma02c,ma04a,ma04b}, the Large Magellanic Cloud \\citep[LMC; e.g.,][]{da06,pop12,grijs13}, M82 \\citep{deg03a,lim13}, NGC 3310 and 6745 \\citep{deg03b}, as well as for stellar population synthesis model comparisons \\citep{deg05,fan12}.  In this paper, we first obtain photometry for all M33 star clusters in our sample (see Section \\ref{s:samp} for definition) based on archival images from the Local Group Galaxies Survey \\citep[LGGS;][]{massey}. Using photometry in the $UBVRI, ugriz$ (Sloan Digital Sky Survey; SDSS) bands and Two Micron All-Sky Survey (2MASS) $JHK$ magnitudes \\citep{2mass}\\footnote{http://www.ipac.caltech.edu/2mass/} when available, the ages and masses of the star clusters in our sample are estimated by comparison of their observed SEDs with updated {\\sc parsec} (version 1.1) isochrones \\citep{bre12}. This paper is organized as follows. Section \\ref{s:data} describes the sample selection and $UBVRI$ photometry. In Section \\ref{s:age} we describe the simple stellar population (SSP) models used, as well as our method to estimate the cluster ages and metallicities. In Section \\ref{s:mass} we present the clusters' mass estimates, and we summarize and conclude the paper in Section \\ref{s:sum}. ",
        "conclusions": "\\label{s:sum}  We have obtained $UBVRI$ photometry for 708 star clusters and cluster candidates in M33, which were selected from \\citet{san10}, including 387, 563, 616, 580, and 478 objects in the $UBVRI$ bands, respectively, of which 276, 405, 430, 457, and 363 did not have previously published $UBVRI$ photometry. The {\\sc SExtractor} code was applied to derive the photometry from LGGS archival images, which cover 0.8 deg$^2$ along the major axis of M33.  We compared our photometry with previous measurements, which showed that our photometry is generally consistent with previous measurements in all filters. The ages, metallicities, and masses of our sample clusters were derived by comparison of their observed SEDs with {\\sc parsec} SSP synthesis models. The fits show that only 205 of the 671 clusters (30.55\\%) are older than 2 Gyr, which is a much smaller fraction than that derived for the M31 globular-like clusters (55.80\\%), suggesting that the M33 cluster population is dominated by young star clusters ($<1$ Gyr). We also note that the mean mass of the M33 star clusters is $1.78\\times10^4M_{\\odot}$, which is much less massive than that of the M31 sample ($2.69\\times10^5M_{\\odot}$) and similar to that of LMC cluster population ($1.51\\times10^4M_{\\odot}$), but higher than that of Milky Way star clusters (including both GCs and OCs: $5.24\\times10^2M_{\\odot}$). This may be related to the fact that the mass of M33 is lower that that of M31, and the gravitational potential is not large enough to produce as many GCs as in M31. Instead, the star-formation history of M33 may be similar to that of the LMC. As for the Milky Way, the recent few Gyr have seen the galaxy undergo quiescent evolution (a low star-formation rate).  On the other hand, the mean metallicity of the M33 clusters ($\\rm [Fe/H]=-1.01$ dex) is much lower than that of the M31 star clusters ($\\rm [Fe/H]=-0.43$ dex) and also of the Milky Way star clusters ($\\rm [Fe/H]=-0.19$ dex), suggesting that its star-formation history has been rather different from those of either M31 or the Milky Way. Based on the cluster mass distributions we also found that the mean mass of star clusters in M33 is similar to that in the LMC but much lower than that in M31 and higher than that in the Milky Way, indicating that M31 underwent more violent star formation than either M33 or the LMC. We also note that stochastic sampling effects can significantly affect our SED fit results for low-mass clusters (i.e., $M_{\\rm cl} \\le 10^3 M_{\\odot}$), potentially leading to large differences in integrated cluster ages, metallicities, and masses \\citep{anders13}. The effects of stochasticity in the clusters' stellar mass functions become weaker with increasing cluster mass. However, even for the highest-mass clusters, $M_{\\rm cl} \\simeq 2 \\times 10^5 M_{\\odot}$, the uncertainties in the derived logarithmic ages are 0.05--0.25 dex; the equivalent uncertainties pertaining to the derived masses are 0.09--0.17 dex. Although these uncertainties are sometimes significant, obtaining accurate, high-resolution spectroscopic observations for statistically large samples of extragalactic star clusters is often prohibitive, particularly if observing time on significantly oversubscribed (large) telescopes is sought. Broad-band imaging and parameter determination based on sophisticated SED fits is the realistic yet a poor man's alternative approach.  The broad-band SED uncertainties can be further reduced for specific age ranges, e.g., for young clusters that exhibit H$\\alpha$ emission, which places additional constraints on the most likely age range. At present, a number of teams are working on quantifying the effects of stochastic sampling; although the underlying message is that such effects may have significant implications in terms of the precision of the derived physical parameters, we argue that a proper understanding of one's uncertainties is of the utmost importance. One should keep these issues in mind when dealing with physical cluster parameters based on broad-band observations."
    },
    "1402/1402.6326_arXiv.txt": {
        "abstract": "Certain configurations of massive structures projected along the line of sight maximize the number of detections of gravitationally lensed $z\\sim10$ galaxies. We characterize such lines of sight with the \\'etendue $\\sigma_\\mu$, the area in the source plane magnified over some threshold $\\mu$. We use the Millennium I and Millennium XXL cosmological simulations to determine the frequency of high $\\sigma_\\mu$ beams on the sky, their properties, and efficient selection criteria. We define the best beams as having $\\sigma_{\\mu>3} >2000$ arcsec$^2$, for a \\zten source plane, and predict $477 \\pm 21$ such beams on the sky. The total mass in the beam and $\\sigma_{\\mu>3}$ are strongly correlated. After controlling for total mass, we find a significant residual correlation between $\\sigma_{\\mu>3}$ and the number of cluster-scale halos ($>10^{14} M_\\sun h^{-1}$) in the beam. Beams with $\\sigma_{\\mu>3} >2000$ arcsec$^2$, which should be best at lensing \\zten galaxies, are ten times more likely to contain multiple cluster-scale halos than a single cluster-scale halo. Beams containing an Abell 1689-like massive cluster halo often have additional structures along the line of sight, including at least one additional cluster-scale ($M_{200}>10^{14}M_\\sun h^{-1}$) halo 28\\% of the time.  Selecting beams with multiple, massive structures will lead to enhanced detection of the most distant and intrinsically faint galaxies. ",
        "introduction": "Accurate and efficient detection of high redshift galaxies is a current limiting problem in observational cosmology. Reionization is expected to occur at a mean redshift of $z \\sim 10.4$ \\citep{Komatsu2011}. Observing how the luminosity function of galaxies changes as reionization progresses is essential in understanding how reionization takes place \\citep{Robertson2013}.  Recent efforts to observe high redshift galaxies have produced over 100 $z\\sim7$ galaxies, and a few \\zten candidate galaxies \\citep{Bouwens2011,Zheng2012,Coe2012a,Ellis2012}. High redshift galaxies are difficult to observe not only due to their distance, but also their intrinsic faintness. Current methods of detecting high redshift galaxies use blank field studies or known strong gravitational lenses.  Using lensing to magnify high redshift source planes has the drawback of reducing the area surveyed per field, though the depth increases. There have been several studies on this tradeoff \\citep{Bouwens2009, Maizy2009}, which conclude that if the effective slope, $\\beta$, of the faint end of the luminosity function $\\phi(L)$ at high redshift, \\begin{equation} \\beta(z) = - \\frac{d(\\log{\\phi})}{d(\\log{L})} \\end{equation} is greater than 1, the gain in depth from lensing overcomes the loss in area. This slope increases with redshift, from $\\sim0.7$ at lower redshifts to $\\sim1$ by $z\\sim7$ \\citep{Bouwens2012,Oesch2012}, though this measurement is still uncertain at the 20\\% level.  Mass along the line of sight (LOS) towards gravitational lenses boosts lensing power \\citep{Wambsganss2005,Hilbert2007,Puchwein2009}. \\citet{Wambsganss2005} find that as the source plane gets further away, a higher fraction of gravitational lenses are composed of multiple significant structures. At $z_{source}\\sim7.5$, about 50\\% of cases studied have a second lens plane with a significant contribution to the total lensing cross section. \\citet{Puchwein2009} claim that LOS matter can boost the strong lensing cross section by up to 50\\%, and that the effect of uncorrelated structure along the line of sight is more significant in increasing the strong lensing cross section than structure in the same lens plane as the cluster. These analyses are primarily aimed at determining how LOS mass affects already good lenses, but raise the possibility  that the best cosmic telescopes could have a form other than a single massive cluster lens.  An analysis of the best configurations of dark matter halos for lensing \\zten galaxies is done by \\citet{Wong2012}. They examine potential lines of sight with mass configurations optimized to maximize the \\'etendue ($\\sigma_\\mu$), the area in the $z\\sim10$ source plane with magnification ($\\mu$) above a certain threshold. Although \\'etendue is typically used in reference to optical systems, we adopt it here to characterize the effective optical system of 3D configurations of gravitational lenses. This provides an advantage over traditional lens characterizations, such as the lensing cross section, which generally refers to the cross-sectional area needed to multiply image source galaxies. For our goal of detecting high redshift galaxies, even modest magnifications provide an advantage, so characterizing beams by \\sm allows us to balance sampling depth and magnified area to find the most effective gravitational lenses. \\citet{Wong2012} consider lines of sight with multiple structures along the line of sight. The division of mass among multiple halos has a large effect on \\smd. Due to interactions between lensing potentials, beams with multiple structures in certain idealized configurations may produce lensing effects comparable to single-cluster beams of twice the total mass. Line of sight lensing beams especially benefit when halos are projected closer on the sky, and if the total mass is split into 2-5 halos of about equal mass.  Knowing the best lens configurations, we can then ask how frequent certain configurations are in the real universe. We want to know the frequency of high \\sm beams, which high \\sm beam configurations are most likely, how often additional significant lensing mass is projected along the line of sight to a massive cluster, and which observables are most successful at finding the best lensing beams.  In this study, we use several new tools to study the properties of the best lensing beams. By defining a lensing \\'etendue, we can better characterize the lenses as observational tools in the search for \\zten galaxies. We frame the effect of LOS mass not as a systematic error to be accounted for, but as a way of boosting the total mass and lensing quality of the beam. The best beams will have very large total masses, including very massive cluster-scale halos.  Drawing halos from simulations requires a large simulation volume in order to sample the high mass tail of the mass function. Here, we use the Millennium XXL simulation \\citep{Angulo2012}. At 3 Gpc $h^{-1}$ on a side, it is much larger than simulations used in previous studies of LOS effects in lensing, producing a larger sample of massive lenses. Using simulation halos rather than toy models provides information about the frequency of halo configurations, which we use to calculate the number of beams that we expect to see on the sky with high \\'etendue.  We focus on maximizing sources detections at $z\\sim10$, as lensing stands to greatly improve number counts and statistics at this redshift. We present the methods used in drawing sample lines of sight from the simulations and calculating their lensing properties in Section 2. We present the results of these analyses and their implications on observing strategy in Section 3. Discussion of some choices in methods is presented in Section 4. We conclude in Section 5. ",
        "conclusions": "We draw massive halos from sample lines of sight in the Millennium I \\citep{Springel2005} and Millennium XXL \\citep{Angulo2012} simulations. We define the quantity \\smd, or \\'etendue, the area in the \\zten source plane where the magnification of the brightest image is greater than a threshold magnification of $\\mu=3$. We characterize those beams with the highest lensing cross-sections \\sm for detecting very high redshift ($z\\sim10$) galaxies.  Our main conclusions are:   \\begin{itemize} \\item We calculate the frequency of high-\\sm beams on the sky, in order to test whether these beams should exist in the real universe. The number of $7^\\prime\\times7^\\prime$ beams with $\\sigma_{\\mu>3} >1000$ arcsec$^2$ out to a redshift of $z=1.4$, and including only halos with $M_{200}>10^{14} M_\\sun h^{-1}$, is $6025 \\pm 78$ over the whole sky. The number of beams with $\\sigma_{\\mu>3} >2000$ arcsec$^2$ is $477 \\pm 21$. These numbers decrease by about 25\\% if MXXL is scaled to a Planck cosmology.   \\item We determine the characteristics of the high \\sm beams. \\sm  increases with total beam mass, and this relation can be fit by a power law with index of $1.36\\pm0.02$.   \\item The highest \\sm beams with $\\sigma_{\\mu>3} >2000$ arcsec$^2$ are ten times more likely to have multiple cluster-scale halos than a single halo. Beams with more halos have on average more total mass, and a resulting higher \\smd. After controlling for total mass, we calculate a significant residual correlation between \\sm and the number of massive halos ($>10^{14} M_\\sun h^{-1}$), and an anti-correlation with the number of lower mass halos ($10^{13} M_\\sun h^{-1}$). In other words, breaking up beam mass into several halos is beneficial to \\sm provided the resulting halos have high cluster-scale masses. This is due to the tradeoff between decreasing lensing power for lower mass halos, and the likelihood of finding two or more close to each other within the specified beam size. These trends hold for different beam sizes, threshold magnifications, and cosmologies.  \\item We study the effect of halo configuration for the special case of two-halo beams. Controlling for mass and redshift separation, we observe a $3\\sigma$ significant anti-correlation of \\sm with the angular separation between the halos. Good \\sm  ($>1000$ arcsec$^2$) beam halos often have halo separations near 2 arcminutes. We observe no significant correlation between \\sm and the redshift offset between the two halos.  \\item Having determined characteristics of the best beams, we identify the observables that are most efficient at  {\\it finding} the best beams. The redshift range that must be covered to find 95 \\% of good \\sm fields ($\\sigma_{\\mu>3} >1000$ arcsec$^2$) is $0<z<1.0$. We construct a classification tree to find regions in observable parameter space that have the highest purity of high \\sm beams. We consider total mass, number of halos, mass of the most massive halo, and minimum and maximum redshifts.  It is most helpful to select beams with total mass above $2\\times10^{15} M_\\sun h^{-1}$, made up of at least 2 structures along the line of sight. Selecting on the mass of the most massive halo or on redshift is not as effective.  \\item We put our analysis of high \\sm beams into the context of other lensing surveys by comparing to known strong lens Abell 1689 and three CLASH \\citep{Postman2012} clusters. For each lens, we select 5 analogue halos in mass and redshift from MXXL, and study possible lines of sight that include that halo. Additional line of sight (LOS) mass causes a systematic increase in \\smd. Beams include at least one additional cluster-scale ($M_{200}>10^{14}M_\\sun h^{-1}$) halo 28\\% of the time. In 20\\% of the total cases, \\sm increases by more than 25\\%, and in 10\\% of the total cases, \\sm increases by more than 80\\%. Excluding LOS mass halos greater than the original lens, \\sm boosts can still be larger than 200\\%. Because of the potential for large changes, and the systematic nature of this effect, LOS mass must be considered in the modeling of these lenses.   \\item The ultimate measure of beam quality is the expected number of \\zten lensed galaxies one would expect to detect. This requires many additional assumptions over \\sm. However, both the number of detections, and the number of detections fainter than $M_{UV}^*=-17.7$, are correlated with \\smd. Beams with \\sm$>2000$ arcsec$^2$ produce 1.1-1.4 times more \\zten detections overall, and 100-250 times more detections fainter than $M_{UV}^*=-17.7$, when compared to a blank field.     \\end{itemize}    Our results will inform searches for high \\sm beams in the real universe, including with LSST \\citep{Ivezic2008}. Our analysis of LRG number density and total LRG luminosity as tracers of massive halos has already shown such beams exist \\citep{Wong2013}. Spectroscopic observations confirm them \\citep{S.MarkAmmonsKennethC.WongCharlesR.Keeton2013}. We have also determined that even single-cluster lenses will be affected by line of sight structure. We argue that these same ``nuisance\" systematics and uncertainties can be used to our advantage in detecting the earliest galaxies."
    },
    "1402/1402.1489_arXiv.txt": {
        "abstract": "We investigate the strength of axisymmetric local pressure maxima (zonal flows) in the outer regions of protoplanetary disks, where ambipolar diffusion reduces turbulent stresses driven by the magnetorotational instability. Using local numerical simulations we show that in the absence of net vertical magnetic fields, the strength of turbulence in the ambipolar dominated region of the disk is low and any zonal flows that are present are weak. For net fields strong enough to yield observed protostellar accretion rates, however, zonal flows with a density amplitude of 10-20\\% are formed. These strengths are comparable to those seen in simulations of ideal MHD disk turbulence. We investigate whether these zonal flows are able to reverse the inward radial drift of solids, leading to prolonged and enhanced concentration as a prelude to planetesimal formation. For commonly assumed mean surface density profiles (surface density $\\Sigma \\propto r^{-1/2}$ or steeper) we find that the predicted perturbations to the background disk profile do not correspond to local pressure maxima. This is a consequence of radial width of the simulated zonal flows, which is larger than was assumed in prior analytic models of particle trapping. These larger scale flows would only trap particles for higher amplitude fluctuations than observed. We conclude that zonal flows are likely to be present in the outer regions of protoplanetary disks and are potentially large enough to be observable, but are unlikely to lead to strong particle trapping. ",
        "introduction": "The radial drift problem in protoplanetary disks results from aerodynamic drag between particles and gas, which orbit at slightly different speeds due to pressure gradients in the gas \\citep{weidenschilling77}. Drift is fastest for particles whose surface area to mass ratio is such that they are marginally coupled, with a dimensionless stopping time $\\tau = m \\Delta v \\Omega^{-1} / |F_{\\rm drag}| \\sim 1$. In the outer regions of protoplanetary disks, at 30 to 100~AU from the star, this condition is met for spherical particles with size $\\sim 1 \\ {\\rm cm}$. Particles of roughly similar size -- including mm-sized particles that can be directly observed in dust continuum observations -- would therefore be expected to drain rapidly from the outer disk, leaving gas behind \\citep{youdin04,hughes12}. In fact, although the extremities of some well-studied disks {\\em are} gas-rich \\citep{andrews12}, mm-sized particles appear to be retained in the outer disk \\citep{ricci10} for substantially longer than is theoretically predicted \\citep{brauer07}. A candidate explanation is that dust is trapped in either permanent or transient local pressure maxima that act to slow radial drift driven by the global pressure gradient \\citep{pinilla12,dittrich13}. If such traps are present, the resulting over-density of solids could locally promote the growth of instabilities leading to planetesimal formation \\citep{chiang10}.  In this paper, we show that under the physical conditions appropriate to the outer regions of protoplanetary disks, zonal flows form spontaneously within magnetohydrodynamic (MHD) disk turbulence. Zonal flows are  transient axisymmetric pressure maxima that were seen in local ideal MHD simulations of accretion  disks \\citep{johansen09}. They are sustained by a geostrophic balance between  pressure gradients and Coriolis forces. Subsequent work has shown zonal flows to be a robust outcome of  ideal MHD disk turbulence in the local limit \\citep{simon12,dittrich13} and has identified similar  structures in global simulations \\cite[e.g.,][]{dzyurkevich10,flock11,uribe11}. What has not been established is  whether, in the presence of the non-ideal MHD effects (Ohmic and ambipolar diffusion, and the  Hall effect) that are important in protoplanetary disks \\citep{armitage11}, zonal flows remain  strong enough to act as efficient traps. Here, we address this question for the outer disk,  where ambipolar diffusion creates an ``ambipolar damping zone\" \\citep{perezbecker11,mohanty13,simon13a,simon13b}.  Our results are based on local (shearing-box) simulations of non-ideal MHD disk turbulence, using methods described in Section 2. In Section 3 we show that within the ambipolar damping zone, the amplitude of predicted zonal flows is a function of the net flux of magnetic field that threads the disk. Zero net flux simulations yield weak turbulence \\citep{simon13a} and very weak zonal flows, while simulations with net fields chosen to yield the observed accretion rates onto young stars \\citep{simon13b} produce prominent zonal flows. In Section 4 we explore whether the derived zonal flows would lead to particle trapping, under the assumption that the perturbed global disk profile can be approximated as the mean global profile multiplied by the local perturbed structure. In Section 5 we compare our results with previous work and discuss the general implications of our results. ",
        "conclusions": "Our primary result is that while we see strong zonal flows, even in the presence of damped turbulence from ambipolar diffusion, the amplitudes obtained from our simulations are not sufficiently large to create pressure maxima where particle trapping can occur.  That the strength of zonal flows in the presence of ambipolar diffusion is similar to that in the ideal MHD limit is surprising. Further examination of the relationship between turbulence levels and zonal flow amplitudes will require a larger parameter study than that performed here.  Our conclusions are less optimistic for the prospects of particle concentration than those of \\cite{pinilla12} who find that particle trapping {\\it can} occur in their model disks. The reason for this difference is the assumed width of the zonal flows, which \\cite{pinilla12} take to be comparatively narrow ($\\delta \\sim 1$). From Fig.~\\ref{eps_delta} it is clear that such a narrow width, if combined with our measured amplitudes, would indeed lead to particle trapping for any reasonable background surface density profile. The problem with this scenario is that our simulations, along with several previous works \\citep{johansen09,simon12}, suggest that $\\delta \\sim 5-10$ is more realistic. With these parameters, trapping particles directly by creating local pressure maxima is significantly more difficult.  There remain several uncertainties in our work.  First, while our simulations include the key MHD physics present in the outer regions of protoplanetary disks (i.e, strong ambipolar diffusion), we have only run a limited number of these simulations,  for a relatively short interval. The roughly 100 orbit duration of our runs is limited by their computational demands,  which remain prohibitive despite the considerable speed up from the use of super-time-stepping.   The bottom panel of Fig.~\\ref{sttx_d_multi1} shows that towards the end of the calculation, the fractional variation in gas density approaches $\\sim 0.3$, and it is possible that integrating this run further would produce sustained zonal flows of roughly this amplitude.  However, it is also possible that these particular zonal flows will decrease in strength and then fluctuate in a stochastic manner, as is observed in IDEAL and previous shearing box simulations \\cite[e.g.,][]{johansen09,simon12}. Second, there are some uncertainties in translating the local simulation results to predictions for trapping in global disk models. High resolution global disk simulations that include ambipolar diffusion, although currently challenging to run, will be essential for a definitive determination of whether outer disk zonal flows can trap particles. Finally, while at face value our results show that particles will not be {\\it trapped}, the pile-up effect due to changes in the radial drift velocity may still play a role in the planet formation process \\citep{johansen06,dittrich13}.  We plan to address these various uncertainties in future work.  Even if zonal flows fail to strongly concentrate particles, their presence may still be observable. At 30~AU, $H / r \\approx 0.1$, and the zonal flow widths found in our simulations correspond to physical scales of the order of 5~AU. Although the perturbations to the gaseous surface density are modest, some degree of enhancement of the particle density due to differences in the radial drift speed across the flow is expected \\citep{johansen06,dittrich13}. Early ALMA results show that there are some disks whose outer regions display dramatic {\\em non-axisymmetric} dust distributions \\citep{vdm13}, and the relatively subtle axisymmetric structures resulting from our zonal flows would likely not be detectable in such systems. It may, however, be possible to observationally constrain the existence and properties of zonal flows by focusing on those disks with smooth radial profiles and the smallest departures from axisymmetry."
    },
    "1402/1402.1676_arXiv.txt": {
        "abstract": "We explore the link between solar energetic particles (SEPs) observed at 1 AU and large-scale disturbances propagating in the solar corona, named after the \\textit{Extreme ultraviolet Imaging Telescope} (EIT) as EIT waves, which trace the lateral expansion of a coronal mass ejection (CME). A comprehensive search for SOHO/EIT waves was carried out for 179 SEP events during Solar Cycle 23 (1997\\,--\\,2006). 87$\\,$\\% of the SEP events were found to be accompanied by EIT waves. In order to test if the EIT waves play a role in the SEP acceleration, we compared their extrapolated arrival time at the footpoint of the Parker spiral with the particle onset in the 26 eastern SEP events that had no direct magnetic connection to the Earth. We find that the onset of proton events was generally consistent with this scenario. However, in a number of cases the first near-relativistic electrons were detected too early. Furthermore, the electrons had in general only weakly anisotropic pitch-angle distributions. This poses a problem for the idea that the SEPs were accelerated by the EIT wave or in any other spatially confined region in the low corona. The presence of weak electron anisotropies in SEP events from the eastern hemisphere suggests that transport processes in interplanetary space, including cross-field diffusion, play a role in giving the SEPs access to a broad range of helio-longitudes. ",
        "introduction": "% \\label{S-Introduction}  Solar energetic particle (SEP) events are transient enhancements of the intensities of electrons above some keV and ions and protons above some hundreds of keV, observed by particle detectors onboard spacecraft or on the ground. The two main candidates for the particle accelerator in the corona are a flare-related reconnection process and a coronal mass ejection (CME) driven shock. Irrespective of the acceleration agent, the particles need to be injected onto magnetic field lines connected to the spacecraft in order to be observed. When the parent activity is in the West, both flares (via rapidly diverging field lines in the low corona: \\opencite{2008A&A...486..589K}) and CMEs (via their large-scale shock front: \\opencite{2006SSRv..123..127S}) can be connected to the magnetic-field lines that reach Earth. This is more difficult to achieve with increasing distance from the optimal magnetic connection to Earth near 60$^\\circ$ western longitude, as predicted by a Parker spiral (PS) model.  Nonetheless, SEP events also occur with activity in the eastern hemisphere. Their origin was discussed in terms of: coronal propagation, diffusive or not, from an accelerator in the associated flare \\cite{1991PCS....21..243K}; acceleration at a CME shock in the interplanetary (IP) space \\cite{1988JGR....93.9555C,1991PCS....21..243K,2006SSRv..123..217K}; and/or propagation in a transient magnetic-field configuration when the Earth is connected to the Sun by an interplanetary coronal mass ejection (ICME) from a previous eruptive solar event \\cite{1991JGR....96.7853R,2012A&A...538A..32M}. In the present work, we investigate whether large-scale traveling disturbances in the low corona, so-called EIT waves, play a role in accelerating SEPs or injecting them into the IP space.  These coronal disturbances were detected for the first time by the \\textit{Extreme ultraviolet Imaging Telescope} instrument (EIT: \\opencite{1995SoPh..162..291D}) onboard the \\textit{Solar and Heliospheric Observatory} (SOHO). On running-difference images, EIT waves \\cite{1998GeoRL..25.2465T,1999ApJ...517L.151T}, also called in the past coronal Moreton waves and recently extreme ultra-violet (EUV) waves, appear as bright fronts that emanate from the active region \\cite{2012SoPh..281..187P}. In the quiet corona they may propagate up to nearly 360 degrees \\cite{2012ApJ...756..143O}, whereas at the boundaries of different magnetic topologies the disturbance undergoes reflections and refractions \\cite{2008ApJ...681L.113V,2009ApJ...691L.123G,2013SoPh..286..201K}. Recent observations from the \\textit{Solar Terrestrial Relations Observatory} (STEREO) and \\textit{Solar Dynamics Observatory} (SDO) spacecraft showed that the EUV structure is closely related to the flanks of the associated CME during its lateral over-expansion \\cite{2010ApJ...708.1639M,2011SoPh..273..421T}.  The nature of this coronal disturbance is subject to an ongoing debate (see reviews by \\opencite{2008SoPh..253..215V}; \\opencite{2009SSRv..149..325W}; \\opencite{2010AdSpR..45..527W}; \\opencite{2011JASTP..73.1096Z}; \\opencite{2012SoPh..281..187P}). There are two opposite interpretations of the observed EUV-intensity fronts. One group interprets them as true fast-mode MHD waves or shocks propagating through the corona and compressing the plasma at its front \\cite{1998GeoRL..25.2465T,2001ApJ...560L.105W}. The other group of models considers the bright EUV front as a consequence of the large-scale restructuring of the coronal magnetic field during the CME lift-off \\cite{1999SoPh..190..107D,2002ApJ...572L..99C,2007ApJ...656L.101A}. A reconciling interpretation, {\\it i.e.} a hybrid of the previous two, was also proposed \\cite{2004A&A...427..705Z}. In the present study, we will use the term ``EIT disturbance\u00b4\u00b4 or ``EIT wave\u00b4\u00b4 in order to refer to the observational phenomenon without any implication on the actual nature of the disturbance.  The idea of a physical relation between the EIT disturbances and SEP events was already investigated in the past with different authors reaching different conclusions for electrons and protons. In an early study, \\inlinecite{1997ESASP.415..207B} compared the propagation of an EIT wave from an eastern active region and the onset of an electron event at SOHO. The EIT wave was just entering the western solar hemisphere when the electron event started, and was therefore not the likely accelerator of the first electrons seen by SOHO. Negative results for electrons were obtained also by \\inlinecite{1999ApJ...519..864K}. They performed a detailed velocity-dispersion analysis of near-relativistic electron events and estimated the initial solar release time for each. For 18 of 58 events, EIT disturbances were found in the SOHO/EIT data. In the six events investigated in more detail, the EIT wave speeds were found to be a factor of two too slow to explain the electron release times at the longitude of best magnetic connection from the Sun to the spacecraft. The authors argued that a 3D disturbance could have higher speed at higher altitude (1.5 solar radii) than an EIT wave that is assumed to propagate at low coronal heights \\cite{2009SoPh..259...49P,2009ApJ...703L.118K}, and thus might be the particle accelerator. The majority of the events were at western heliospheric longitudes.  The onsets of protons up to 25~MeV were, however, compatible with acceleration as the EIT wave reached the vicinity of the IP field line connected to the spacecraft. This conclusion was reached in the case study of the 24 September 1997 proton event by \\inlinecite{1999ApJ...510..460T}, where the initial injection of the $\\gtrsim$10 MeV protons was found to be during the period when an eastern EIT wave was traversing the western hemisphere of the Sun. In a recent, very detailed, multi-spacecraft analysis \\inlinecite{2012ApJ...752...44R} show that the arrival time of the CME flank and the associated EIT wave at the respective footpoints of the spacecraft-connected Parker spirals could explain the arrival timing of the first energetic electrons and protons at different spacecraft. \\inlinecite{2009ApJ...704..469M} argued that both electrons and protons observed in December 2006 by {\\it Ulysses} and \\textit{Advanced Composition Explorer} (ACE), in association with an active region near the eastern limb, could be attributed to EIT disturbances reaching magnetic field lines connected to the spacecraft.  In summary, the question of whether EIT waves contribute to SEP acceleration has so far no unique answer. The present work addresses the subject by combining a statistical assessment with a timing analysis, where the arrival of EIT waves at the root of the Earth-connected PS in the corona is compared with the onset time of the SEP events near Earth. Starting with a comprehensive list of particle events (Section~\\ref{S-All_events_data}) in Solar Cycle 23 (1997\\,--\\,2006), we performed an extensive search in the SOHO/EIT data for associated large-scale coronal disturbances by visually identifying the wave fronts. Results are summarized in a catalog of EIT waves.  In order to test if an EIT wave accelerates SEPs far from the parent active region, we focus (Section~\\ref{S-Eastern}) on events related to eruptive activity in the eastern solar hemisphere. We excluded those events where the SEPs were detected while the spacecraft was in or near an ICME, which could provide a transient direct magnetic-field connection to the eastern solar hemisphere. The onset of the SEP events -- deka-MeV protons and near relativistic electrons observed by \\textit{Electron Proton Helium Instrument} (EPHIN) part of the \\textit{Comprehensive Suprathermal and Energetic Particle Analyzer} (COSTEP) onboard SOHO -- and the onset and speed of the EIT waves were determined, together with the expected arrival of the EIT wave at the PS footpoint. By comparing this arrival time and the SEP onset one can decide if the EIT waves propagate fast enough to accelerate the particles that arrived first at the spacecraft, Section~\\ref{S-Results}. Additional information on the electrons is provided from the observations of the anisotropy characteristics from \\textit{Electron, Proton, and Alpha Monitor} (EPAM) onboard ACE. We summarize the results in Section~\\ref{S-Summary} and discuss the impact of our findings on scenarios of the SEP origin from eastern activity in Section~\\ref{S-Disc}. ",
        "conclusions": "\\label{S-Summary}  We conducted a survey of the association between SEP events and EIT waves in the Solar Cycle 23. For a more detailed investigation of the timing relationship between the EIT wave and the SEP event onset, we then focused on eastern solar events. Whenever possible we used SOHO/EIT imaging observations to extrapolate the arrival time of the EIT wave at the foot of the Earth-connected PS field line in the western hemisphere and tested the consistency with the onset time of the SEP events, both near-relativistic electrons and deka-MeV protons, at the SOHO spacecraft. The observational results are the following: \\begin{enumerate}[i)] \\item The large majority of SEP events are associated with EIT waves (87$\\,$\\%). The rate is comparable for events associated with eastern and western activity with an uncertainty of some percent due to the uncertain association of solar activity with some of the SEP events. The converse is not true: only a minority of EIT waves (about 10$\\,$\\%) are accompanied by SEP events. \\item Eastern SEP events are observed within ICMEs (17$\\,$\\%), in the standard solar wind (69$\\,$\\%) or in a situation where an ICME is within a day from the Earth and can have caused a major perturbation of the medium in which the SEPs travel. The latter category is ambiguously defined. But the fraction of events detected within ICMEs is consistent with the earlier findings of \\inlinecite{1991JGR....96.7853R} of the order of 15$\\,$\\%. \\end{enumerate}  A more detailed study of the timing of SEPs (onset delay with respect to the parent solar activity and slope of the initial rise of the time profiles, characterized by the rise time $[t_{\\rm r}]$) and EIT waves was carried out for the eastern SEP events detected in the standard solar wind:  \\begin{enumerate}[i)] \\item SEP events from eastern eruptive activity that are observed in the standard solar wind tend to have a broader range of rise times and onset delays with respect to the parent eruptive activity than events observed in ICMEs. \\item The onset delays of electrons and protons are correlated. About half of the SEP events detected in the solar wind show no clear velocity-dispersed onset, as seen by a comparison of deka-MeV protons and near-relativistic electrons. In the others, the temporal profile of the proton intensities rises after the electrons. \\item The time constants of the early rise of the intensity--time profiles of electrons and protons are also correlated. \\item Both for electrons and protons, the onset delays are correlated with the rise times. \\item The onset delays -- and hence also the early rise times -- are not ordered by the connection distance, that is the longitudinal difference between the erupting active region and the spacecraft-connected PS. \\item The arrival of the first deka-MeV protons is consistent with the scenario of an acceleration in the low corona as the EIT wave reaches the footpoint of the nominal PS to the spacecraft. But in the majority of cases the first electrons are detected earlier than expected by the scenario. \\item With one exception, the pitch-angle distributions of near-relativistic electrons are weakly anisotropic or isotropic. This is different from events observed within ICMEs and electron events associated with activity in the western hemisphere. \\end{enumerate}"
    },
    "1402/1402.5263_arXiv.txt": {
        "abstract": "{Hot subdwarf B stars (sdBs) are core helium-burning stars located on the extreme horizontal branch. About half of the known sdB stars are found in close binaries. Their short orbital periods of 1.2 h to a few days suggest that they are post common-envelope systems. Eclipsing hot subdwarf binaries are rare, but important to determine the fundamental stellar parameters. Low-mass companions are identified by the reflection effect. In most cases the companion is a main sequence star near the stellar mass limit.\\\\ Here we report the discovery of an eclipsing hot subdwarf binary SDSS J162256.66+473051.1 (J1622) of very short orbital period (0.0697 d), found in the course of the MUCHFUSS project. The lightcurve shows grazing eclipses and a prominent reflection effect. An analysis of the light- and radial velocity (RV) curves indicated a mass ratio of $q=$ 0.1325, an RV semiamplitude $K=47.2\\rm\\,km\\,s^{-1}$, and an inclination of $i=72.33^\\circ$. We show that a companion mass of 0.064 $M_{\\rm \\odot}$, well below the hydrogen-burning limit, is the most plausible solution, which implies a mass close to the canonical mass (0.47 $M_{\\rm \\odot}$) of the sdB star. Therefore, the companion is a brown dwarf, which not only survived the engulfment by the red-giant envelope, but also triggered its ejection, and enabled the sdB star to form.\\\\ The rotation of J1622 is expected to be tidally locked to the orbit. However, J1622 rotates too slowly ($v_{\\rm rot}=74.5\\pm 7\\rm\\,km\\,s^{-1}$) to be  synchronized, challenging tidal interaction models.} ",
        "introduction": "Hot subdwarfs (sdBs) are core helium-burning stars with very thin hydrogen envelopes, that are not able to sustain hydrogen shell burning \\citep{heber:2009}. To form such objects, the progenitor has to lose almost all of its hydrogen envelope. The high percentage of 50\\% of close binaries \\citep{maxted:2001,napi} suggests that these sdBs are formed via binary evolution, which is responsible for the required large mass loss on the red giant branch (RGB). Such sdB binaries are formed via a common envelope phase or by stable Roche lobe overflow. The other half of sdB stars, however, appears to be single. Hence, binary evolution seems to be irrelevant at first glance. Nevertheless, is it possible to form a single star through merging of the components of a binary star. To form a helium core burning object like an sdB star, requires two helium white dwarfs to merge and ignite helium burning. The binary components are driven into a merger by gravitational wave radiation \\citep{webbink,it}. The merged object is expected to rotate rapidly. However, rotation velocities of single sdB stars are very slow \\citep[$\\rm <=10\\,km/s$,][]{geier:2012}, which is at variance with the merger scenario. Therefore, the merger scenario may explain a few cases of exceptionally fast rotators \\citep{geier:2011_3,geier:2013_rot}. Moreover, binary population synthesis predicts a wide mass distribution around 0.52 $M_{\\rm \\odot}$ for the sdB, resulting from the merger channel \\citep{han:2002,han:2003}. The empirical mass distribution derived by asteroseismology \\citep{fontaine,vangrootel:epjwc}, however, shows a sharp peak at $0.47 M_\\odot$ for single sdB stars, which is inconsistent with the prediction of the merger scenario.  An alternative scenario proposes the engulfment and possible destruction of a substellar object within a common envelope (CE) as formation channel for single sdB stars \\citep{soker,nt}. Close binary sdB stars with very low-mass stellar companions are known for decades \\citep{kilkenny:1978,menzies:1986}. When eclipsing, such systems are named HW Vir stars. Most of the dozen known HW Vir systems have periods as short as 0.1 days. Hence, they must have undergone a common envelope and spiral-in phase. Binary population synthesis shows that sdBs in such post-common envelope systems should have masses around $0.47\\,M_{\\odot}$, the so-called canonical mass \\citep{han:2002,han:2003}. The companions in HW Vir type binaries have masses close to, but usually exceeding the nuclear-burning limit of $\\simeq0.08\\,M_{\\odot}$ and are, therefore, very late M-dwarf stars.  Whether a substellar companion would also be able to contribute enough energy and angular momentum to unbind a common envelope, has been under debate. The discovery of a brown dwarf orbiting an sdB star in the HW Vir type binary SDSS J082053.53+000843.4 demonstrated that substellar companions are able to form an sdB \\citep{geier}. The subsequent discoveries of two Earth-sized bodies orbiting a single pulsating sdB, which might be the remnants of one or two more massive companion evaporated in the CE-phase \\citep{charpinet:nature,bear}, provides further evidence that substellar companions play an important role in the formation of close binary and single sdB stars alike. One of the open questions is now, how massive the companion has to be, to trigger the mass loss of the red giant to form an sdB star.  Here we report on the discovery of a short period eclipsing hot subdwarf with a substellar companion in the course of the MUCHFUSS project. The project Massive Unseen Companions to Hot Faint Underluminous Stars from SDSS (MUCHFUSS) aims at finding sdBs with compact companions, such as massive white dwarfs  (M  > 1.0 $M_{\\odot}$), neutron stars, or black holes. Details about target selection and the scope of the project can be found in \\citet{geier:2011_2}. As one of two selection criteria, sdBs with radial velocity (RV) variations on timescales of half an hour or less were selected from the Sloan Digital Sky Survey (SDSS). These targets are either sdBs with compact, massive companions and comparatively long orbital periods, or short-period binaries with low-mass companions like the system described here.  To distinguish low-mass stellar or substellar companions from white dwarfs, and other more compact objects, we also started a photometric follow-up. As HW Vir systems have characteristic lightcurves, this turned out to be the ideal approach to find such binaries. A low-mass, cool companion can be identified due to the so-called reflection effect, which results from the large difference in temperature of both components. The companion is heated up on one side, and, therefore, a sinusoidal variation of the flux with the phase can be observed, as the heated hemisphere emits more flux. This effect is also visible at rather low inclinations, where no eclipses are observed. If eclipses are present, a combined spectroscopic and photometric analysis allows us to put tight constraints on the binary parameters. In the course of the MUCHFUSS project we already discovered three new HW Vir systems \\citep{geier,tuc_schaff} and a pulsating sdB in a reflection effect binary \\citep{ostenson:2013}.  In this paper we present the observations and analysis of a HW Vir star, discovered by the MUCHFUSS project. We describe the spectroscopic and photometric observations in Sect. 2 and their analysis in Sect. 3 (spectroscopy) and 4 (photometry). Evidence for the brown dwarf nature of the companion is given in Sect. 5 and the lack of synchronisation of the sdB star is discussed in Sect. 6. Finally, we conclude and present suggestions, how to improve the mass determination. ",
        "conclusions": "We performed a spectroscopic and photometric analysis of the eclipsing hot subdwarf binary J1622, that was found in the course of the MUCHFUSS project. The atmospheric parameters of the primary are typical for an sdB star in a HW Vir system. Mass-radius relations were derived for both components. The mass of the sdB star is constrained from spectroscopy (surface gravity) to 0.28 $M_{\\odot}$ to 0.64\\,$M_{\\odot}$. The mass of the companion can be constrained by theoretical mass-radius relations to lie between 0.055\\,$M_{\\rm \\odot}$ and 0.075\\,$M_{\\rm \\odot}$, which implies that the sdB mass is between 0.39\\,$M_{\\rm \\odot}$ and 0.63\\,$M_{\\rm \\odot}$.  Assuming small corrections of about  5-10\\% to the radius due to the inflation of the companion by the strong irradiation from the primary, a companion mass of  0.064\\,$M_{\\rm \\odot}$ appears to be the most plausible choice resulting in a mass of the sdB close to the canonical mass of $0.47 M_{\\rm\\odot}$ star. Accordingly the companion is a substellar object. This is the second time a brown dwarf is found as a close companion to an sdB star. J1622 provides further evidence that substellar objects are able to eject a common envelope and form an sdB star. Finding more such systems, will help to constrain theoretical models \\citep{soker,nt}.  An important result of the spectral analysis is that the sdB star is rotating slower than expected, if its rotation was locked to its orbit, as is observed for the very similar system PG\\,1017$-$036. The non-synchronous rotation of J1622 is at variance with the predictions of tidal interaction models unless the sdB star is very young. The ratio of the rotational to the orbital period is close to a 2 to 3 resonance. However, it has to be investigated further, if such a configuration can survive the common envelope phase.  Future investigations should aim at the detection of spectral lines from the secondary, in order to measure its radial velocity, that is to turn the system into an  double-lined spectroscopic binary, which would allow to pin down the mass ratio. Emission lines of the companion's irradiated atmosphere have been detected in the sdOB system AA\\,Dor \\citep{aador} and most recently also for the prototype HW\\,Vir \\citep{tuc_maya}.  Since J1622 is quite compact and the irradiation of the companion strong, such emission lines might be detectable in high-resolution spectra of sufficient quality. Accurate photometry is needed to confirm or disprove any asymmetries in its lightcurve hinting at flows at the surface of the companion."
    },
    "1402/1402.2675_arXiv.txt": {
        "abstract": "We try to explain the subpulse drift phenomena adopting the space-charge limited flow (SCLF) model and comparing the plasma drift velocity in the inner region of pulsar magnetospheres with the observed velocity of drifting subpulses. We apply the approach described in a recent paper of \\cite{lt12}, where it was shown that the standard estimation of the subpulse drift velocity through the total value of the scalar potential drop in the inner gap gives inaccurate results, while the exact expression relating the drift velocity to the gradient of the scalar potential should be used instead. After considering a selected sample of sources taken from the catalog of \\cite{wes06} with coherently drifting subpulses and reasonably known observing geometry, we show that their subpulse drift velocities would correspond to the drift of the plasma located very close or above the pair formation front. Moreover, a detailed analysis of PSR B0826-34 and PSR B0818-41 reveals that the variation of the subpulse separation with the pulse longitude can be successfully explained by the dependence of the plasma drift velocity on the angular coordinates. ",
        "introduction": "The research on pulsar magnetospheres started 45 years ago from the pioneering paper of \\cite{gj69}, where it was shown that the magnetic field, together with the fast rotation of the pulsar, generates strong electric fields tending to pull out charged particles from the surface of the neutron star and to form the plasma magnetosphere. However, the very next arising question - how many particles will actually leave the surface of the star under the action of this force - is still a subject of scientific debate. In fact, the magnetic field which generates the pulling out electric field also leads to the substantial increase in the cohesive energy of the surface charged particles (positive ions in larger degree than electrons), making the outer layer of the star very dense and strongly bound. \\cite{ml07} showed that for each magnetic field intensity there exists a critical surface temperature, above which the particles are able to freely escape from the surface of the star. The magnetic field of pulsars is typically inferred from observations under the assumption that it is strictly dipolar, thus providing $B_d = 2\\times 10^{12}\\sqrt{P\\dot{P}\\times10^{15}}\\mathrm{G}$ (here $P$ is the period and $\\dot{P}$ is the period derivative of the pulsar). Using this formula, it turns out that the majority of pulsars satisfy the condition for the free particle outflow. However, in many works, starting from \\cite{rs75} it was suggested that the magnetic field close to the surface of the pulsar should have a multipole structure, with the surface magnetic field several orders of magnitude larger than the estimated $B_d$. This idea is supported by X-ray observations \\citep{zsp05,kpg06,pkwg09}, and several studies have been performed searching for a mechanism of generation and maintenance of such small-scale strong magnetic fields [see \\cite{ggm13} and references therein].  The amount of charged particles extracted from the surface of the star by the rotationally induced electric field is a key aspect of any pulsar magnetosphere model. The model of \\cite{rs75}, for example, assumes that no particle leaves the pulsar surface and there is a vacuum gap formed above the star with a huge difference in the scalar potential between the bottom and the top, \\ie $\\sim 10^{12} \\mathrm{V}$. According to this idea, the gap will be periodically discharged and rebuilt, making it intrinsically \"non-stationary\". On the contrary, in the \"stationary\" space-charge limited flow (SCLF) model of \\cite{as79} there is a free flow of charged particles from the pulsar surface. The partially screened gap model of \\cite{gmg03} combines some features of both the previous models and is based on  the assumption that the vacuum gap discharge will lead to the back flow of charged particles bombarding and heating the surface of the star, causing the thermal ejection of ions from the surface and partially screening the original accelerating electric field.  The choice among different models of pulsar magnetosphere should be made by accurate comparison of their predictions with the results of observations. A very interesting phenomena serving as a diagnostic tool is the subpulse drift. Pulsar radio emission comes to us in the form of pulses, which may look very different among each other and typically consist of individual subpulses. Although the average pulse profile is very stable and represents a unique \\emph{fingerprint} of each pulsar, a range of successive pulses plotted on top of each other in a so-called pulse stack, quite often shows an organized phase shift of the subpulses forming drift bands. This phenomenon was reported for the first time in \\cite{dc68}, while the first systematic analysis of ``drifting'' pulsars dates back to \\citet{b81} and \\citet{r86}. So far, the largest statistical study of the phenomenon has been presented in the works of \\cite{wes06} and \\cite{wse07}, who considered a sample of 187 pulsars, $55\\%$ of which show drifting subpulses. Usually the subpulse drift bands are characterized by the horizontal separation between them in subpulse longitude, $P_2$, and the vertical separation in pulse periods, $P_3$. The subpulse behavior of the individual pulsars may be rather complex and demonstrate smooth or abrupt change of drift direction, phase steps, longitude and frequency dependence of the separation $P_2$, or even presence of subpulses drifting in the opposite directions at the same time. {For the graphical representation of the periods $P_2$ and $P_3$ we refer the reader to Fig.~1 of \\cite{wes06}, where one can find several examples of the different subpulse behavior of individual sources.}  The vacuum gap model, and especially the partially screened gap model, has been widely used to explain the subpulse drift phenomena \\citep{gs00,mgp00,gmg03,ggks04,bggs07,getal08}, while for a long time the SCLF model has been regarded unable to account for it. However, recent analytical \\citep{lt12} and numerical \\citep{t10b,ta13} progresses have shown that the door can be left open even for the SCLF model. The main goal of this paper is to explain the subpulse drift velocity in the framework of the relativistic SCLF model, without addressing the question of the generation mechanism of the plasma features responsible for the appearance of the subpulses, while trying to compare the results with the available observational data.  The plan of the paper is the following. In Sect.~\\ref{oldmodels} we give a brief review of the models used so far to explain the subpulse phenomena and motivate our choice to concentrate on the SCLF model. In Sect.~\\ref{velocitySCLF} we present the basic equations to explain the subpulse drift velocity in the framework of the SCLF model. In Sect.~\\ref{alt} we consider a set of pulsars from the catalog of \\cite{wes06} with coherently drifting subpulses, and try to deduce in which regions of the pulsar magnetosphere the SCLF model would predict the plasma with the observed velocities. In Sect.~\\ref{individualpulsars} we focus on two specific sources, PSR B0826-34 and PSR B0818-41, trying to account for their phenomenology. Finally, Sect.~\\ref{conclusions} is devoted to the summary of the results obtained and to the conclusions. ",
        "conclusions": "  \\begin{itemize} \\item the SCLF model predicts the subpulse drift velocities compatible to the observed ones at heights above the surface of the star close to the pair formation front;  \\item the angular dependence of the plasma drift velocity in the SCLF model provides a natural explanation for the variation of the subpulse separation $P_2$ along the pulse. In particular it may explain the curved subpulse driftbands of PSR B0818-41 and the range of the observed subpulse velocities of PSR B0826-34.  \\end{itemize}  These results suggest that the role of the SCLF model in explaining the drifting subpulse phenomena has been underestimated, calling for additional investigations and systematic comparisons with all available observations."
    },
    "1402/1402.2443_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  In the present Universe, the Hubble parameter $H^2$ is so small in the unit of the Planck mass $M_P^2 = 1/ G$: \\begin{equation} H^2/M_P^2 = 10^{-120}. \\label{problem}\\end{equation} Its energy scale may be related to that of neutrino mass as \\begin{equation} H^2M_P^2\\sim m_\\nu^4. \\end{equation} It has been long thought to be exactly zero before its recent observations. The present Universe is well described as the space-time with a positive cosmological constant, namely de Sitter (dS) space. However we have not understood why the Hubble parameter (Dark Energy) is so small.  It is likely that we cannot explain its magnitude in the present theoretical framework: relativistic quantum field theory. As is well known, the quantum field theory in the flat space-time is intimately connected to critical phenomena in equilibrium physics. In fact the space-time with a positive cosmological constant is not stationary. The quantum field theory in dS space may belong to nonequilibrium physics. In fact we need to employ the Schwinger-Keldysh formalism instead of the familiar Feynman-Dyson formalism.  In particular, the existence of a cosmological event horizon in dS space may give rise to interesting quantum effects. It generates the fluctuations of the space-time metric with a scale invariant spectrum. Inflation theory explains the origin of large scale structures of space-time by such a mechanism. We may need to take the fluctuations at the super-horizon scale seriously as they may eventually come back into the sub-horizon scale.  In the Poincar\\'{e} coordinate, the metric in dS space is \\begin{align} ds^2=-dt^2+a^2(t)d{\\bf x}^2,\\hspace{1em}a(t)=e^{Ht}, \\end{align} where the dimension of dS space is taken as $D=4$ and $H$ is the Hubble constant. In the conformally flat coordinate, \\begin{align} (g_{\\mu\\nu})_\\text{dS}=a^2(\\tau)\\eta_{\\mu\\nu},\\hspace{1em}a(\\tau)=-\\frac{1}{H\\tau}. \\end{align} Here the conformal time $\\tau\\ (-\\infty <\\tau < 0)$ is related to the cosmic time $t$ as $\\tau=-\\frac{1}{H}e^{-Ht}$. We assume that dS space begins at an initial time $t_i$ with a finite spatial extension. After a sufficient exponential expansion, dS space is well described locally by the above metric irrespective of the spatial topology.  The metric is invariant under the scaling transformation \\begin{align} \\tau\\to C\\tau,\\hspace{1em}x^i\\to Cx^i. \\label{scaling}\\end{align} It is a part of the $SO(1,4)$ dS symmetry. The central issue is whether the dS symmetry could be broken due to IR effects. Note that this is a large gauge (coordinate) transformation. As explained shortly, we may need to introduce an IR cut-off in order to regulate the IR divergences of the propagators of the minimally coupled modes. The IR cut-off breaks the scale invariance explicitly. Therefore the equivalent question is to ask whether there exist IR divergences in physical observables.  Let us consider free propagators of a massless and minimally coupled scalar field $\\varphi$ and a massless conformally coupled scalar field $\\phi$ \\begin{align} \\langle\\varphi(x)\\varphi(x')\\rangle&=\\frac{H^2}{4\\pi^2}\\big\\{\\frac{1}{y}-\\frac{1}{2}\\log y+\\frac{1}{2}\\log a(\\tau)a(\\tau')+1-\\gamma\\big\\}, \\label{minimally0}\\end{align} \\begin{align} \\langle\\phi(x)\\phi(x')\\rangle&=\\frac{H^2}{4\\pi^2}\\frac{1}{y}, \\label{conformally0}\\end{align} where $\\gamma$ is Euler's constant and $y$ is the dS invariant distance \\begin{align} y=\\frac{-(\\tau-\\tau')^2+({\\bf x}-{\\bf x}')^2}{\\tau\\tau'}. \\end{align} It should be noted that the propagator for a massless and minimally coupled scalar field has the dS symmetry breaking logarithmic term: $\\log a(\\tau)a(\\tau')$.  To explain what causes the dS symmetry breaking, we recall the wave function for a massless and minimally coupled field \\begin{align} \\varphi_{\\bf p}(x)=\\frac{H\\tau}{\\sqrt{2p}}(1-i\\frac{1}{p\\tau})e^{-ip\\tau+i{\\bf p}\\cdot{\\bf x}}. \\end{align} Well inside the cosmological horizon where the physical momentum $P\\equiv p/a(\\tau)\\gg H \\Leftrightarrow p|\\tau|\\gg 1$, this wave function approaches to that in Minkowski space up to a cosmic scale factor \\begin{align} \\varphi_{\\bf p}(x)\\sim\\frac{H\\tau}{\\sqrt{2p}}e^{-ip\\tau+i{\\bf p}\\cdot{\\bf x}}. \\end{align} On the other hand, the behavior outside the cosmological horizon $P\\ll H$ is \\begin{align} \\varphi_{\\bf p}(x)\\sim\\frac{H}{\\sqrt{2p^3}}e^{i{\\bf p}\\cdot{\\bf x}}. \\end{align} The IR behavior indicates that the corresponding propagator has a scale invariant spectrum. As a direct consequence of it, the propagator has a logarithmic divergence from the IR contributions in the infinite volume limit.  To regularize the IR divergence, we introduce an IR cut-off $\\varepsilon_0$ which fixes the minimum value of the comoving momentum \\begin{align} \\int^H_{\\varepsilon_0a^{-1}(\\tau)} dP. \\end{align} With this prescription, more degrees of freedom come out of the cosmological horizon with cosmic evolution. Due to the increase, the propagator acquires the growing time dependence which spoils the dS symmetry \\cite{Vilenkin1982,Linde1982,Starobinsky1982}. In tribute to its origin, we call the dS symmetry breaking term the IR logarithm. Physically speaking, $1/\\varepsilon_0$ is recognized as an initial size of universe when the exponential expanding starts. For simplicity, we set $\\varepsilon_0=H$ in (\\ref{minimally0}). We recall here that the physical momentum $P=p/a(\\tau)$ is invariant under the scaling transformation (\\ref{scaling}). Our fundamental hypothesis is to adopt the dS invariant UV cut-off $P<\\Lambda_\\text{UV}$. In this prescription UV contributions from $H<P<\\Lambda_\\text{UV}$ is time independent. In contrast, the IR contributions from $H/a(\\tau)<P<H$ could grow with cosmic expansion.  Actually the metric fluctuations always contain a scale invariant spectrum just like the massless and minimally coupled scalar field. In dealing with the quantum fluctuation of the metric whose background is dS space, we adopt the following parametrization: \\begin{align} g_{\\mu\\nu}=\\Omega^2(x)\\tilde{g}_{\\mu\\nu},\\ \\Omega(x)=a(\\tau)e^{\\kappa w(x)}, \\label{para1}\\end{align} \\begin{align} \\det \\tilde{g}_{\\mu\\nu}=-1,\\ \\tilde{g}_{\\mu\\nu}=(e^{\\kappa h(x)})_{\\mu\\nu}, \\label{para2}\\end{align} where $\\kappa$ is defined by the Newton's constant $G$ as $\\kappa^2=16\\pi G$.  The Lagrangian of Einstein gravity on the $4$-dimensional dS background is \\begin{align} \\mathcal{L}_\\text{gravity}&=\\frac{1}{\\kappa^2}\\sqrt{-g}\\big[R-6H^2\\big]\\label{gravity}\\\\ &=\\frac{1}{\\kappa^2}\\big[\\Omega^2\\tilde{R} +6\\tilde{g}^{\\mu\\nu}\\partial_\\mu\\Omega\\partial_\\nu\\Omega-6H^2\\Omega^4\\big], \\notag \\end{align} where $\\tilde{R}$ denotes the Ricci scalar constructed from $\\tilde{g}_{\\mu\\nu}$. In order to fix the gauge with respect to general coordinate invariance, we adopt the following gauge fixing term \\cite{Tsamis1992}: \\begin{align} \\mathcal{L}_\\text{GF}&=-\\frac{1}{2}a^{2}F_\\mu F^\\mu, \\label{GF}\\\\ F_\\mu&=\\partial_\\rho h_\\mu^{\\ \\rho}-2\\partial_\\mu w+2h_\\mu^{\\ \\rho}\\partial_\\rho\\log a+4w\\partial_\\mu\\log a. \\notag \\end{align} In this paper, the Lorentz indexes are raised and lowered by the flat metric $\\eta^{\\mu\\nu}$ and $\\eta_{\\mu\\nu}$ respectively.  After decomposing the spatial part of the metric and diagonalizing the quadratic action: \\begin{align} h^{ij}=\\tilde{h}^{ij}+\\frac{1}{3}h^{kk}\\delta^{ij}=\\tilde{h}^{ij}+\\frac{1}{3}h^{00}\\delta^{ij}, \\end{align} \\begin{align} X=2\\sqrt{3}w-\\frac{1}{\\sqrt{3}}h^{00},\\hspace{1em}Y=h^{00}-2w, \\label{diagonalize}\\end{align} we find that some modes of gravity behave as the massless and minimally coupled scalar field and the other modes behaves as the massless and conformally coupled mode: \\begin{align} \\langle X(x)X(x')\\rangle&=-\\langle\\varphi(x)\\varphi(x')\\rangle, \\label{minimally}\\\\ \\langle\\tilde{h}^i_{\\ j}(x)\\tilde{h}^k_{\\ l}(x')\\rangle&=(\\delta^{ik}\\delta_{jl}+\\delta^i_{\\ l}\\delta_j^{\\ k}-\\frac{2}{3}\\delta^i_{\\ j}\\delta^k_{\\ l})\\langle\\varphi(x)\\varphi(x')\\rangle, \\notag\\\\ \\langle b^i(x)\\bar{b}^j(x')\\rangle&=\\delta^{ij}\\langle\\varphi(x)\\varphi(x')\\rangle, \\notag \\end{align} \\begin{align} \\langle h^{0i}(x)h^{0j}(x')\\rangle&=-\\delta^{ij}\\langle\\phi(x)\\phi(x')\\rangle, \\label{conformally}\\\\ \\langle Y(x)Y(x')\\rangle&=\\langle\\phi(x)\\phi(x')\\rangle, \\notag\\\\ \\langle b^0(x)\\bar{b}^0(x')\\rangle&=-\\langle\\phi(x)\\phi(x')\\rangle, \\notag \\end{align} where $b$, $\\bar{b}$ denote the ghost and anti-ghost fields.  Since we focus on the dS symmetry breaking effects, we may introduce an approximation. We can neglect the conformally coupled modes of gravity (\\ref{conformally}) since they do not induce the IR logarithm. In such an approximation, the following identity holds \\begin{align} h^{00}\\simeq 2w\\simeq\\frac{\\sqrt{3}}{2}X. \\label{diagonalize1}\\end{align} Our aim is to investigate quantum IR effects from the minimally coupled modes of gravity (\\ref{minimally}) on microscopic matter dynamics. To do so, we derive the effective equation of motion which takes account of  the quantum effects due to soft gravitons. Before investigating specific models, we review the general procedure to derive the effective equation of motion in the next section. ",
        "conclusions": "\\setcounter{equation}{0}  The dS symmetry may be broken due to growing IR effects. Since there exist massless minimally coupled modes in the metric fluctuations, such an effect may be a ubiquitous feature of quantum gravity in dS space. In particular we have found that the couplings acquire time dependence and evolve logarithmically with the scale factor of the Universe. We thus propose that the relative evolution speeds of the couplings are the physical observables in dS space.  We observe these effects can be explained by the nontrivial conformal mode dependence of the operators. We recall here the analogous renormalization takes place in 2d quantum gravity. There the scaling dimensions of the local operators and equivalently couplings to them are renormalized by the scale invariant quantum fluctuation of the metric. We have recalled that the cosmological constant is screened by IR quantum fluctuations in the semiclassical regime. We believe therefore analogous renormalization takes place in 4d dS space due to the scale invariant quantum IR fluctuations of the metric. We argue that the IR logarithmic effects originate from the conformal mode dependence in the effective propagators of the scalar and Dirac fields. In fact we have shown such a quantum effect is the inevitable consequence of the general covariance. Thus we have made a strong case for the validity of our claim when these fields are involved.  The magnitude of this effect is of $O(H^2/M_P^2)$. Since this factor is very small $(10^{-120})$ now, it seems impossible to detect such an effect at first sight. However this factor may be compensated by UV divergences in the case of the cosmological constant which is quartically divergent with respect to the UV cut-off. If so, such an effect could lead to very significant consequences: \\begin{itemize} \\item {Dark energy may decay with a small but finite $n < 10^{-3}$.} \\item Inflation may split SUSY in the early universe. \\end{itemize}  Let us recall BPHZ renormalization procedure. We assume that the theory is renormalized at the $n$-th order with appropriate counter terms. In the next order, the new divergences occur as over-all divergences since sub-divergences are assumed to be subtracted. If the couplings are time independent, we can renormalize these new local divergences by covariant counter terms. In this way renormalizability implies the absence of UV-IR mixing effect and vise versa. On the other hand, we cannot do so if the couplings become time dependent. The UV-IR mixing effect could arise as the incomplete cancellation of the divergences due to time evolution of the couplings.   So UV-IR mixing effects could occur in non-equilibrium and non-renormalizable field theory. Although our reasoning is different, we agree with Polyakov that the cosmological constant problem may very well be of this nature \\cite{Polyakov}.  Let us consider the density fluctuation in the cosmic microwave background (CMB) due to scalar modes. The standard argument relates it to the two point function of the minimally coupled scalar field (inflaton) at the horizon crossing as it is identified with the curvature perturbation which remains constant after the horizon crossing in the comoving gauge \\cite{Mukhanov1990,Maldacena2002}. We may be able to estimate the scalar and tensor perturbations from the sub-horizon theory side by following it up to the horizon crossing. We have shown that the leading IR effects can be absorbed into the renormalization of the fields and couplings inside the cosmological horizon. We expect that this picture still holds when the space time is approximately dS such as in the inflation era. We hence believe that our results have an implication for the super-horizon soft graviton loop effect on CMB. If so, it shows that the leading IR effects on the scalar two point functions due to soft gravitons are absent as they can be renormalized away. The time evolution of the couplings modifies slow roll parameters $\\epsilon$ and $\\eta$ in general.  The curvature and tensor perturbations are shown to stay constant after the horizon crossing in the comoving gauge at the tree level in a single field inflation. It is because they approach pure gauge. The nontrivial potential for them are forbidden due to general covariance. The issues on the higher loop corrections are reviewed in \\cite{Tanaka}. The IR logarithmic corrections influence the CMB spectrum by making the microscopic parameters of the inflation theory time dependent. We believe that it is important to understand them when we investigate microscopic physics behind inflation.  We illustrate one of such a possibility. The self-tuning effect of the cosmological constant due to UV-IR mixing could alter the classical slow roll picture. The slow roll parameters are \\begin{align} \\epsilon =\\frac{1}{16\\pi}(\\frac{M_PV'}{V})^2\\sim 4\\pi\\frac{\\dot{\\phi}^2}{H^2}\\frac{1}{M_P^2},\\hspace{1em}\\eta=\\frac{1}{8\\pi}\\frac{M_p^2V''}{V}. \\end{align} We note that the slow roll parameter $\\epsilon$ acts like $n/2$ in (\\ref{index}) \\begin{align} H^2 \\sim a(t)^{-2\\epsilon}. \\end{align} So the quantum effects could significantly alter the classical slow roll picture when $\\epsilon$ becomes as small as $n$. In the case of the single field inflation, $n$ may be small as $\\lambda_4$ is strongly constrained by the observation. In a generic case, $n$ could be as large as $10^{-3}$ since it could arise from different fields rather than the inflaton. The quantum effect alters the classical slow roll picture in general if the tensor-to-scalar ratio $r$ is as small as $0.01$. Such a parameter region will be explored soon."
    },
    "1402/1402.2505_arXiv.txt": {
        "abstract": "Pulsars, if existing and detectable in the immediate vicinity of the massive black hole (MBH) in the Galactic center (GC), may be used as a superb tool to probe both the environment and the metric of the central MBH. The recent discovery of a magnetized pulsar in the GC suggests that many more pulsars should exist near the MBH.  In this paper, we estimate the number and the orbital distribution of pulsars in the vicinity of the MBH in the GC by assuming that the pulsar progenitors, similar to the GC S-stars, were captured to orbits tightly bound to the MBH through the tidal breakup of stellar binaries. We use the current observations on both the GC S-stars and the hypervelocity stars to calibrate the injection rate(s) of and the dynamical model(s) for the stellar binaries.  By including the relaxation processes, supernova kicks, and gravitational wave radiation in our simulations, we estimate that $\\sim97-190$ ($9-14$) pulsars may presently orbit the central MBH with semimajor axes $\\leq4000\\AU$ ($\\leq1000\\AU$), which is compatible with the current observational constraints on the number of the GC pulsars. The semimajor axis and the pericenter distance of the pulsar closest to the central MBH are probably in the range of $\\sim120-460\\AU$ and $\\sim2-230\\AU$, respectively. Future telescopes, such as the SKA, may be able to detect a significant number of pulsars with semimajor axis smaller than a few thousand AU in the GC. Long-term monitoring of these pulsars would be helpful in constraining both the environment and the metric of the central MBH. Our preferred model also results in about ten hyperfast pulsars with velocity $\\ga1500\\kms$ moving away from the Milky Way. ",
        "introduction": "Long term monitoring of the motion of the S-stars\\footnote{Hereafter denoted as the GC S-stars.} located at the Galactic center (GC) demonstrates exclusively the existence of a massive black hole (MBH) in the GC \\citep[e.g.,][]{Ghezetal08, Gillessenetal09}. Although the source confusion in observations has so far inhibited the discovery of any star rotating around the MBH with an orbital period $<10$~yr \\citep{Meyer12}, young stars and compact stars, including neutron stars and stellar mass black holes, are expected to exist in this region \\citep[e.g.,][]{MG00, Angelil10, liu12, ZLY13}. If some of these stars are radio pulsars, which would be promising in observational detection and providing a superb clean tool to probe the gravitational field of the central MBH \\citep[e.g.,][]{PT79, Pfahl04, Angelil10}. It becomes especially encouraging that a magnetar, a highly magnetized pulsar, is recently discovered at a parsec (pc) distance from the central MBH. This magnetar has been used to reveal the magnetic field in the GC \\citep[see][]{Rea13, Eatough13, Kennea13, Mori13}. Magnetars are a rare type of pulsars, and the discovery of a magnetar in the GC suggests that many more normal pulsars exist there.  The GC S-stars may be the captured components of stellar binaries that were tidally broken up in the vicinity of the central MBH as proposed by \\citet[][see also \\citealt{GL06, Lockmann08, Perets09, AM13, ZLY13}]{Gould03}.  The other components of those broken-up binaries may be ejected to the Galactic halo as hypervelocity stars \\citep[HVSs; e.g.,][]{Hills88, YT03, Brown05, Edelmann05, Hirsch05, Brown12}.  \\citet[][hereafter ZLY13; see also \\citealt{LZY10}, and \\citealt{ZLY10}]{ZLY13} find that the current observations on both the GC S-stars and the HVSs are compatible with the model that their progenitorial stellar binaries are originated from the young stellar disk(s) at a distance of $0.04-0.5\\pc$ from the MBH \\citep{LuJ09,Bartko09}.  Assuming that both the HVSs and the GC S-stars are produced by the tidal breakup mechanism, the injection rate of stellar binaries to the immediate vicinity of the central MBH can be constrained by using current observations on the numbers and distributions of these two populations.  Some of the captured/ejected massive components (with mass $\\geq 9\\msun$; \\citealt{Heger03}) of the broken-up binaries may collapse and explode to form pulsars at the end of their main-sequence lives. Therefore, the number of the pulsars currently existing in the GC can be reasonably estimated if adopting the tidal breakup mechanism to produce the massive GC S-stars.  \\citet{Pfahl04} estimated the number of the pulsars existing in the vicinity of the central MBH ($\\la 4000\\AU$)\\footnote{Different units for the distance of a star from the GC MBH are chosen in different references, e.g., pc in \\citet{liu12} and \\citet{Gillessenetal09}, mas ($10^{-3}\\,$arcsec) in \\citet{Ghezetal08}, and gravitational radius $r_{\\rm g}$ in \\citet{Angelil10}, respectively. For the GC MBH system, we note here that $1 \\text{mas} \\simeq 3.88\\times 10^{-5}\\,\\text{pc} (R_{\\rm GC}/8 \\text{kpc}) \\simeq 8\\text{AU} (R_{\\rm GC}/8 \\text{kpc}) \\simeq 202 r_{\\rm g} (M_{\\bullet}/4\\times 10^6\\msun)^{-1} (R_{\\rm GC}/8\\text{kpc})$. In this paper, the MBH mass is set to be $M_{\\bullet} = 4\\times 10^6\\msun$, the distance from the GC to the sun is $R_{\\rm GC} = 8 \\text{kpc}$, and we mainly use the physical distance unit AU to describe the distance of a GC star from the central MBH.} to be $\\sim 1000$ (see also a new estimate on the pulsar generation rate in the GC by \\citealt{DO13}). In this paper, we revisit the estimation by including more detailed and realistic considerations of various physical processes involved in, such as the dynamical evolution of the pulsars and their progenitors, the kicks received by the pulsars at their birth, and the orbital decay due to the gravitational wave (GW) radiation, etc.  This paper is organized as follows. In Section~\\ref{sec:simu}, we perform a large number of full three-body experiments to realize the tidal breakup processes of stellar binaries in the vicinity of the central MBH by adopting different binary injection models. We follow the dynamical evolution of the captured components, by considering the resonant and the non-resonant relaxations due to background stars, and the decay due to the GW radiation.  If a captured (or an ejected) star ends its main-sequence life and explodes to form a pulsar before the end of the simulation, the motion of the pulsar is assumed to receive a kick at its birth, based on a kick-velocity distribution obtained from observations. We set the end of the simulations as the present time. We follow the kinematic motions of the ejected components in the Galactic potential and obtain the number and the  spatial distribution of HVSs at the present time. The parameters describing the initial settings for each hypothesized injection model can be calibrated by using current observations as done in ZLY13. With the calibrated model parameters, the number and the spatial distribution of these pulsars are estimated in Section~\\ref{sec:result}. We further estimate the probability distributions of the semimajor axis and the pericenter distance of the inner most pulsar. Discussion and conclusions are given in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we estimate the number and orbital distribution of the pulsars currently existing in the vicinity of the MBH in the GC, by assuming that these pulsars are formed from the captured components of the massive binaries tidally broken up in the vicinity of the MBH.  We follow the orbital evolution of those pulsars (and its progenitorial stars) by considering the effects of the limited stellar lifetime, the supernova kicks, the resonant and the non-resonant relaxations due to background stars, and the GW radiation. For the model that can simultaneously re-produce the statistical properties of both the observed GC S-stars and HVSs (i.e., the Disk-IM2 model), the number of pulsars with semimajor axis $\\leq 4000\\AU$ (or $\\leq 1000\\AU$) is estimated to be about $97-190$ (or $9-14$). Among them, the closest one to the central MBH probably has a semimajor axis and pericenter distance in the range of $116-461\\AU$ and $2-228\\AU$, respectively.  The existence of pulsars in the GC is suggested by a number of indicative observations, including the pulsar wind nebular discovered by \\citet{Wang06}, many X-ray binaries discovered by \\citet{Muno05} and \\citet{Bower05} at a distance within one parsec from the MBH, and a magnetar recently discovered within a parsec distance from the central MBH \\citep[see][]{Rea13, Kennea13, Mori13, Eatough13}, etc. The discovery of the magnetar is most encouraging because magnetars, i.e., the magnetized pulsars, are a rare type of pulsars; the discovery of a single magnetar in the central pc region is almost exclusively indicating the existence of a substantial number of normal pulsars within that region.  Recent systematic searches for pulsars did not detect any but gave an upper limit of $\\sim 90$ normal pulsars in orbits within the central pc of Sgr~A* \\citep[e.g.,][]{Macquart10, Eatough12}.  The non-detection of pulsars in the GC for those searches may be partly due to that pulsar signals are strongly smeared by the hyper-strong scattering of radio waves by the turbulent ionized gas within the central $100\\pc$ of Sgr\\,A*. The pulse broadening has a strong frequency dependence, $\\propto \\nu^{-4}$, making it nearly impossible to detect any pulsar in the GC at the typical observing frequencies of $<1.4$GHz. It may be possible to avoid of this smearing problem by searching pulsars at high frequencies ($>10$GHz). \\citet{Chennamangalam13} revisited the constraint on the number of potential observable pulsars according to the current observations by adopting more realistic assumptions; they found a conservative upper limit $\\sim 950$ for the pulsars in orbits within the central pc.  The non-detection of pulsars in the GC is still compatible with our estimate of the number of the pulsars there ($\\sim 97-190$, see Table~\\ref{tab:t1}). According to the method to estimate the number of observable pulsars developed in \\citet[][]{Codres97}, the fraction of observable pulsars in the GC by the current facilities is $\\sim 10^{-2}$ \\citep[see Equation~2 in][]{Pfahl04}. As the expected total number of pulsars in the GC, with semimajor axes $\\leq 4000\\AU$, is only $\\sim 97-190$, the detectable pulsars would be only $\\sim 1-2$ (or even zero) by the current facilities.  Future radio telescopes, such as the SKA, may be able to detect $10\\%$ of the pulsars \\citep[see Equation~2 in][]{Pfahl04}. We expect that ten to twenty (about one) pulsars will be discovered with semimajor axis $\\leq 4000\\AU$ ($\\leq 1000\\AU$) in the GC by the SKA.  Among the ``detected'' pulsars, the closest one to the central MBH probably has semimajor axis and pericenter distance in the range of $468-1423\\AU$ and $21 - 791\\AU$, respectively. For other models that can produce $17$ GC S-stars as observed, the total number of the ``detected'' pulsars, with semimajor axes $\\leq 4000\\AU$ ( $\\leq 1000\\AU$), are about $4-50$ ($1-10$), and the one closest to the central MBH may have a semimajor axis smaller than that resulting from the Disk-IM2 model. Long term monitoring of those pulsars close to the central MBH would be helpful in constraining the environment and the metric of the central MBH.  Our preferred model also results in about ten hyperfast pulsars with velocities $\\geq 1500\\kms$ moving away from the Milky Way, and most of them are produced by supernovae explosions of the massive HVSs ejected from the GC. Some of these hyperfast pulsars may be detected by the SKA in the future. However, it is not easy to confirm that these hyperfast pulsars are originated from the GC solely by their moving directions."
    },
    "1402/1402.4818_arXiv.txt": {
        "abstract": "Many exoplanets have been discovered with radii of 1$-$4 R$_\\earth$, between that of Earth and Neptune. A number of these are known to have densities consistent with solid compositions, while others are ``sub-Neptunes'' likely to have significant H$_2$-He envelopes. Future surveys will no doubt significantly expand these populations. In order to understand how the measured masses and radii of such planets can inform their structures and compositions, we construct models both for solid layered planets and for planets with solid cores and gaseous envelopes, exploring a range of core masses, H$_2$-He envelope masses, and associated envelope entropies.  For planets in the super-Earth/sub-Neptune regime for which both radius and mass are measured, we estimate how each is partitioned into a solid core and gaseous envelope, associating a specific core mass and envelope mass with a given exoplanet. We perform this decomposition for both ``Earth-like'' rock-iron cores and pure ice cores, and find that the necessary gaseous envelope masses for this important sub-class of exoplanets must range very widely from zero to many Earth masses, even for a given core mass.  This result bears importantly on exoplanet formation and envelope evaporation processes. ",
        "introduction": "\\label{intro}  \\begin{figure}[htp] \\includegraphics[bb = 18 144 330 718,clip,width=\\columnwidth]{exoplanets.eps} \\caption{Comparison of known exoplanets and Solar System planets with our models of simple Fe core/MgSiO\\(_3\\) mantle planets, pure water (Ice VII) planets, and planets with H$_2$-He envelopes. Earth-like is defined as 32.5\\% core mass fraction (CMF) and Mercury-like is defined as 70\\% CMF \\citep{2007ApJ...669.1279S}. The range of possible mass/radius values for several planets lie squarely within the area occupied by Fe core/MgSiO\\(_3\\) mantle planets, making them excellent terrestrial exoplanet candidates. A number of others lie between the pure MgSiO$_3$ and pure Ice VII curves, making them, perhaps, candidates for ``water worlds,'' or for possessing small H$_2$-He envelopes. Another population of planets at a wide range of masses is consistent only with deeper H$_2$-He envelopes. The measured mass and radius values and references for the plotted planets are given in Table \\ref{planet_list}.} \\label{exoplanets} \\end{figure}  The detection of thousands of candidate exoplanets with a wide range of masses and radii motivates the study of the general structure of\tplanetary bodies. While early detection methods heavily favored large planets with masses and radii near those of Jupiter, the recent trend has been toward lower masses and radii, some of which appear to be terrestrial, e.g. Kepler-10b, which was recently confirmed as a planet with radius \\(1.416^{+0.033}_{-0.036}R_\\earth\\), mass \\(4.56^{+1.17}_{-1.29}M_\\earth\\), and average density $8.8^{+2.1}_{-2.9}$ g cm\\(^{-3}\\) \\citep{2011ApJ...729...27B}. For comparison, the average densities of the Earth and Venus are \\(5.5\\) g cm\\(^{-3}\\) and \\(5.2\\) g cm\\(^{-3}\\), respectively.  \\citet{2013arXiv1311.0329L} have suggested that it is likely that planets larger than about 1.75 R$_\\earth$ (based on mass-radius relations) have hydrogen/helium envelopes that contribute significantly to their radii.  In particular, they find that a planet's radius alone provides a first-order estimate of its composition, specifically, the H$_2$-He mass fraction. There is some uncertainty in this limit. For example, \\citet{2014ApJ...783L...6W} adopt a maximum solid planet radius of 1.5 R$_\\earth$, based in a maximum in the density distribution at $\\sim$1.5 R$_\\earth$ and $\\sim$7.6 M$_\\earth$ and \\citet{Marcy.in.press} interpret this as a transition radius of 2.0 R$_\\earth$, given an observed decrease in density from 1.5 to 2.0 R$_\\earth$.  Similarly, \\citet{2011ApJ...727...86R} suggests that envelope accretion onto a core, leading to a significant gaseous envelope, begins at a core mass of 10 M$_\\earth$, or perhaps larger if the planets form close to their stars. However, recent observations of known exoplanets suggest that envelope accretion begins, on average, at a lower mass (as found by \\citealt{2014ApJ...783L...6W}), and some individual planets appear to acquire gaseous envelopes at very low masses. For example, Kepler-51b has been measured to have a mass of $2.1_{-0.8}^{+1.5}$ M$_\\earth$ and a radius of $7.1\\pm 0.3$ R$_\\earth$, corresponding to a density of 0.03$^{+0.02}_{-0.01}$ g cm$^{-3}$ \\citep{2014arXiv1401.2885M}, clearly indicating a mostly gaseous composition.  Recent space-based missions such as {\\it Kepler} \\citep{2010Sci...327..977B} and {\\it CoRoT} \\citep{2006ESASP1306...33B} had photometric precision capable of measuring transits by Earth-sized planets. In the first 16 months of the {\\it Kepler} Mission, 207 Earth-sized (\\(R_p<1.25 R_\\earth\\)) and 680 super-Earth-sized (\\(1.25 R_\\earth < R_p < 2 R_\\earth\\)) planetary candidates were reported \\citep{2013ApJS..204...24B}, suggesting a large number of solid planet candidates given a $\\sim$1.75 R$_\\earth$ cutoff. Figure \\ref{exoplanets} provides a comparison of planets with measured radii and masses with theoretical mass-radius curves that we have generated for various simple planet compositions, including pure iron, Earth-like, Mercury-like, and pure silicate.\\footnote{We define ``Earth-like'' to refer to a composition of 32.5\\% Fe and 67.5\\% MgSiO$_3$, and, when specifying a planet with a significant volatile mass fraction, a 13:27 ratio of Fe to MgSiO$_3$ mass. Similarly, we define ``Mercury-like'' to refer to a composition of 70\\% Fe and 30\\% MgSiO$_3$, and a 7:3 ratio if a significant volatile mass fraction is specified \\citep{2007ApJ...669.1279S}. Note, however, that recent observations from MESSENGER suggest that Mercury's iron content may be closer to 73\\% \\citep{2013JGRE..118.1204H}.} We also include on Figure \\ref{exoplanets} a pure water$-$in the form of Ice VII$-$mass-radius curve and three curves for models with gaseous envelopes: an ``Earth-like'' solid core and H$_2$-He mass fractions of 0.1\\%, 1\\%, and 10\\%.\\footnote{Throughout our paper, we use an envelope composition of 75\\% H$_2$ and 25\\% He by mass.}  \\begin{figure}[htp] \\includegraphics[width=\\columnwidth]{planets.density.eps} \\caption{Known extrasolar planets plotted against constant density curves, including Earth's density (green) and Neptune's density (blue). Densities ($\\rho$) are given in g cm$^{-3}$.} \\label{planets_density} \\end{figure}  As Figure \\ref{exoplanets} demonstrates, the mass/radius values of several known planets are consistent with an iron/rocky composition, and consistent with an Earth-like composition in particular.\\footnote{One planet on the plots (Kepler-131c) has a measured density too high to be consistent with even a pure iron composition, and two others (Kepler-68c and Kepler-406c) have densities that are consistent only with the pure iron curves, likely due to the difficulty of making accurate mass measurements for smaller planets.} On the other hand, a number of observed planets have densities between those of pure water and pure silicate and are consistent with models of both ``water-worlds'' with a high water content and no significant envelopes, and models with H$_2$-He envelopes. Based on our models, the radius of a pure water, 10-M$_\\earth$ planet is 2.5 R$_\\earth$, which is near an observed break in the planet occurrence function in {\\it Kepler} observations \\citep{2012ApJS..201...15H,2013ApJ...778...53D}. However, it is not known which of these models (if either) is dominant in this radius range or whether this break reflects an actual difference in composition between planets smaller and larger than 2.5 R$_\\earth$.  In Figure \\ref{exoplanets}, we note a very wide spread in the mass-radius distribution for planets more massive than $\\sim$2 M$_\\earth$, with a variation of $\\sim$2 R$_\\earth$ at a given mass. For planets $\\lesssim$8 M$_\\earth$, this range overlaps with an Earth-like composition with no significant gaseous envelope. For higher-mass planets, some of these are also consistent with a no-envelope model if they have a sufficient water content, but we also observe planets with large radii that are consistent only with a structure that includes a deep H$_2$-He envelope, even at low masses ($\\gtrsim$2 M$_\\earth$).  In Figure \\ref{planets_density}, we plot the extrasolar planets against constant density curves. Earth-density (green) and Neptune-density (blue) curves are included. The figure shows that density can vary by a factor of $\\sim$5 between individual planets with the same radius. This large scatter makes it very difficult to fit any precise trends in radius with increasing mass.  In order to determine whether a given mass/radius pair indicates a solid composition, however, solid planet models and models of planets with gaseous envelopes must be constructed and compared with the data. These exoplanets may potentially have a wide range of possible compositions and temperatures and a similarly wide range of possible gaseous envelopes, so models must be able to be adjusted accordingly. Solid exoplanet models are important in both cases, since they may be used to model solid cores of planets with gaseous envelopes by applying a non-zero-pressure boundary condition at the core-envelope interface.  While there is a rich history of exoplanet structural modeling, there are a number of important areas which have yet to be investigated. For planets with gaseous envelopes, the effects of irradiation and atmospheric heating are poorly understood, and the degeneracies of envelope mass, envelope entropy, and core mass have not been explored in detail. Moreover, the implications of the large scatter in the mass-radius distribution, particularly on the search for Earth-like planets, are only beginning to be addressed.  For solid planets, equations of state for planetary materials at the pressures found in planets are subject to a degree of uncertainty, and different EOS models produce different results, which warrant analysis of the uncertainty in the modeled mass and radius values.  Our paper investigates the uncertainties and degeneracies in exoplanet modeling, particularly of planets with H$_2$-He envelopes, in order to gain a better understanding of what is measurable in observed exoplanets. We compute mass-radius curves over the range of 0.1 to 20 M$_\\earth$ for a variety of planet attributes, and explore the possibility of determining a precise core-envelope decomposition from mass and radius observations. We study planets with both ``Earth-like'' cores and ice cores, which may both be of interest depending on whether planets with gaseous envelopes form beyond the snow line. We also compute mass-radius curves for various compositional profiles for solid planets. Note that observations suggest that planets more massive than $\\sim 4-8$ M$_\\earth$ are likely to possess significant H$_2$-He envelopes \\citep{2014ApJ...783L...6W}, and the same is expected to be true of planets with radii larger than 1.5$-$2.0 R$_\\earth$ \\citep{2013arXiv1311.0329L,2014ApJ...783L...6W,Marcy.in.press}, so our results for purely solid planets likely apply only to smaller and less massive objects. Conversely, our results for planets with H$_2$-He envelopes will apply to planets larger than $\\sim 4-8$ M$_\\earth$ or 1.5$-$2.0 R$_\\earth$.  While e.g. \\citep{2013arXiv1311.0329L} perform evolutionary calculations to produce planetary structural models, because of the large uncertainties in the parameters that go into these calculations, we do not do this, and we believe that they do not constitute an improvement over our method. Specifically, the ages of known exoplanets are, in most cases, not measured, so that a wide range of ages is possible. Metallicities affect opacities, which in turn affect the cooling rate. They also affect the mean molecular weight and scale height of the atomsphere. Because the formation mechanism of these planets is not known, it is not appropriate a priori to estimate their uncertainties by analogy with Uranus and Neptune. In addition, irradiation is usually treated in an ad hoc manner by averaging the stellar flux over the entire planet rather than modeling day-to-night heat redistribution, which \\citet{2010ApJ...722..871S} show has a significant effect on the net cooling rate.  All of these effects will influence the outcome of an evolutionary model, resulting in large uncertainties. Therefore, we sidestep these uncertainties by focusing on a single variable, the entropy of the envelope, and compute structural models over a range of entropies. The entropy and surface gravity of a planet completely determine its structure (with a small correction for metallicity), so our models are, in effect, the same as those provided by evolutionary calculations.  Section \\ref{previous} summarizes the previous work modeling solid exoplanets and sub-Neptunes. Section \\ref{hydrogen} presents our models with H$_2$-He envelopes and the effect on radius of varying the core mass, envelope mass, and entropy in the envelope. In Section \\ref{fits}, we explore the core-envelope decomposition and produce a fit to the mass-radius relation of \\citet{2014ApJ...783L...6W}. Section \\ref{error} gives an overview of the quantitative effects of varying our model parameters. Section \\ref{internal} demonstrates our code's output with density-pressure profiles of planets, central pressures and densities, and envelope base pressures. Section \\ref{solid} presents our models of solid planets, and Section \\ref{conclusions} summarizes our overall conclusions. Our modeling procedure, associated code, and our studies and selection of our equations of state are described in the Appendix. ",
        "conclusions": "\\label{conclusions}  We have investigated a range of exoplanet models for various core masses, gaseous envelope masses, and envelope entropies, and compared them with mass and radius observations. Some of our representative modeled planet properties are tabulated in Tables \\ref{0.1menv_6.0_table}-\\ref{Fe_planet_table} for a variety of planet compositions. We have explored models with both ``Earth-like'' rock-iron cores and ice cores to account for the possibility of the formation of planets with gaseous envelopes in both the warm and cold regions of their ``solar'' systems, and we have investigated the correction for the presence of a radiative atmosphere. We also considered varying silicate and water fractions for solid planets.  We have decomposed observed exoplanets into core mass and envelope mass components for both rock-iron cores and ice cores. Based on measured masses and radii, we find that the envelope mass, $M_{env}$, may vary over a wide range of values from zero to tens of percent of the total mass over a wide range of total masses, except that for the higher-mass ``sub-Neptune'' (or ``mini-Neptune'') planets, a nonzero envelope mass is always required to fit a rock-iron core. Thus, planetary formation and evaporation models need to account for the very wide range of core masses and envelope masses derived.  In general, an ice core model requires a smaller envelope mass (and larger core mass) to fit the same mass-radius pair as an Earth-like core model, because of the larger radius of the core. Entropy also has a systematic effect on the core-envelope decomposition: an envelope with higher entropy is hotter and has a lower density, so that a smaller envelope mass is needed to fit the same mass-radius pair. The correction for the presence of a radiative atmosphere is also significant, reducing the envelope mass needed to fit the data by 30\\%-50\\% for large envelopes, and sometimes eliminating the need for a convective envelope entirely in the case of small atmosphere masses. Therefore, given these uncertainties, we have derived a range of possible core and envelope masses for known Neptune- and sub-Neptune-sized exoplanets.  While a few planets have large envelope mass fractions of $\\sim 22\\%-35\\%$ (when corrected for the radiative atmosphere), as shown in Figure \\ref{planets_h2o_fits} and Table \\ref{planet_list}, most of the ``sub-Neptune'' planets that have been observed are dominated by their core masses (Figures \\ref{planets_fits} and \\ref{planets_h2o_fits}); that is, the envelope comprises only a small fraction of the total mass. At the same time, for core masses of $M_c\\gtrsim 5$ M$_\\earth$, the planetary radius is very sensitive to the envelope mass and also to the entropy, so that the observed radius can serve as a proxy for the properties of the envelope, subject to the degeneracies of envelope mass, entropy, and the depth of the radiative atmosphere.  For solid planets up to 20 R$_\\earth$, we find that only planets with radii $\\lesssim 1.5-2.0$ R$_\\earth$ can be purely terrestrial (iron core plus silicate mantle). Observationally, the largest planet that is consistent with our terrestrial models is Kepler-20b, with $M=8.7\\pm 2.2$ M$_\\earth$ and $R=1.91_{-0.21}^{+0.12}$ R$_\\earth$. In a transitional regime of $\\sim$1.75-2.8 R$_\\earth$ (which may correspond with an observed break in the planetary occurence function, \\citealt{2013ApJS..204...24B}), the abundance of both water and H$_2$-He may be important, but the effect of the water fraction on the radius diminishes for larger planets.  Determining the composition and structure of solid planets from mass and radius observations is more ambiguous than for planets with gaseous envelopes due to significant degeneracies and uncertainties. There is overlap in radius with different iron fractions and water fractions over the mass range at which planets with potentially-solid compositions have been observed. Further research is needed to standardize equations of state for planetary models and reduce uncertainties.  There remains a degeneracy between composition and envelope entropy, and more detailed atmospheric and evolutionary models are needed to estimate the atmospheric entropy of exoplanets (to a precision of $\\sim0.1 k_B/B$) in order to make more accurate determinations of their compositions and envelope/core mass partitions. On the observational side, more precise mass measurements are needed to better constrain core masses, as well as envelope masses in the case of smaller planets. More planet detections with overall better statistics are also needed to fully populate the mass-radius diagram and determine the distinct populations of planets (if any) in this regime."
    },
    "1402/1402.1110_arXiv.txt": {
        "abstract": "{The spin of a number of black holes (BHs) in X-ray binaries (XBs) has been predicted (and, in at least three cases, confirmed by observations) by using a binary stellar evolution model with Case-C mass transfer . The rotational energy of such BHs is sufficient to power up (long) gamma-ray bursts and hypernovae (GRBs/HNe) and still leave a Kerr BH behind. However, strong magnetic fields (B fields) and/or dynamo effects in the interior of a BH-progenitor star may be capable of rapidly depleting the angular momentum from the stellar core, hence, preventing the formation of a collapsar.  Thus, even if binaries can produce Kerr BHs, most of their rotation is acquired from accreting the stellar mantle, with a long delay between the formation of the BH and its spin up. Hence, not being good sources of GRBs.  We study the necessary conditions to produce GRBs by the progenitors of such BHs. Tidal-synchronization and Alfv\\'en timescales are compared for B fields of different intensities threading trough He stars. We search  for a B-field range which allows tidal spin up all the way into stellar core but prevents its slow down during differential rotation phases. Energetics for producing a strong B field during core collapse, which allows for GRB central engines, are also estimated. An observationally-reasonable choice of parameters is found. }  \\resumen{Utilizando evoluci\\'on estelar de sistemas binarios (con transferencia de masa Caso C) el esp\\'in de varios agujeros negros (ANs) en binarias de rayos X ha sido predicho, observaciones posteriores han confirmado al menos tres de estas predicciones. La energ\\'ia rotacional de \\'estos ANs es suficiente para producir destellos de rayos gama (GRBs) largos acompa\\~nados por explosiones de hipernova (HN) y a\\'un as\\'i dejar un AN de Kerr. No obstante, campos magn\\'eticos intensos y/o d\\'inamos en el interior de la estrella progenitora del AN pueden extraer toda la reserva de momento angular en el n\\'ucleo estelar previniendo la formaci\\'on de un col\\'apsar. Esto es, a\\'un en binarias capces de formar ANs de Kerr, la mayor parte de la rotaci\\'on del AN puede venir del manto estelar, lo cual produce un largo retraso entre la formaci\\'on del AN y la adquisici\\'on de su esp\\'in. Este escenario descartar\\'ia a estas binarias como posibles precursoras de GRBs.  En este trabajo se realiza un estudio sobre las condiciones necesarias para que las progenitores de ANs de Kerr en binarias de rayos X puedan producir GRBs. Escalas de tiempo de sincronizaci\\'on por fuerzas de marea y de Alfv\\'en son comparadas para campos magn\\'eticos de diferentes intensidades en estrellas de He. Se esfect\\'ua una b\\'usqueda para un rango de campos magn\\'eticos que permitan sincronizaci\\'on por fuerzas de marea hasta el n\\'ucleo estelar pero que prevengan la fuga de momento angular durante las subsiguientes etapas de rotaci\\'on diferencial. Se estima tambi\\'en el costo energ\\'etico de la producci\\'on de un campo magn\\'etico durante el colapso gravitacional que permita la formaci\\'on de un GRB. Finalmente, se encuentra un rango de par\\'ametros congruente con observaciones. }   \\addkeyword{binaries: general} \\addkeyword{black hole physics} \\addkeyword{gamma-ray bursts: general} \\addkeyword{stars: evolution} \\addkeyword{stars: magnetic field} \\addkeyword{stars: massive}  \\begin{document} ",
        "introduction": "\\label{sec:Intro}  Massive stars, those capable of leaving a massive core which may collapse into a black hole (BH) tend to lose a substantial amount of mass through winds. The material leaving the star is at the surface, where most of the angular momentum of the star resides. Thus, wind mass loss implies substantial spin angular momentum loss to these stars. If the magnetic field is considerable at the stellar surface and beyond, the lever arm increases further helping the spin down. Furthermore, evolution into the giant and supergiant branch drasdtically increases the radii and the wind mass loss. Hence, evolved stars lose most of the angular momentum left in the star after the main sequence. This picture is enhanced by metallicity, as wind-mass loss is increases along with it.  Long gamma-ray bursts (LGRBs) are thought to be the product of a massive, rapidly-rotating, gravitationally-core-collapsing (GCC) star. During GCC said stars collapse first from the poles as no centrifugal force prevents the material from falling. This clears the way for jets to propagate along the rotational axis of the star as a lower density region, and lower resistance path is created. It is also likely that during GCC differential rotation ensues and any fossil magnetic field becomes entwined forming magnetic towers with their axis coinciding with the rotational one. Therefore, further helping the jet stay focused to drill its way out of the star.  In \\citep{2002ApJ...575..996L} a binary model is proposed where angular momentum is injected back into the stellar core at the very last hour. This is done with the idea that little angular momentum will be lost at that stage even if large wind-mass-loss rates prevail. Plus the orbit will still be a large reservoir of angular momentum for the spin (considering an orbit which does not become too wide from the little mass loss that may still occur). For binaries with Case C mass transfer (after He core burning), a common envelope phase will follow. This allows for the orbital separation to shrink at the expense of removing the H outer shell of the star. As the companion star aproaches the He core of the supergiant, the density increases and, eventually, the later is again confined to its shrunken Roche lobe, thus, ending the common envelope phase. The core of the supergiant, now, likely a Wolf-Rayet (WR) star, still fills a substantial portion of its Roche lobe nonetheless. From the tidal-synchronization timescales estimated in \\citet{1975A&A....41..329Z} and \\citet{1977A&A....57..383Z}, which depend on a high power (6 to 8.5) of the ratio $(a/R)$ ($a$ being the orbital separation and $R$ the stellar radius) it can be shown \\citep[see, e.g.,][]{2007Ap&SS.311..177V} that the massive star will synchronize with the orbit within its remaining lifetime before GCC.  Using this model, \\citet{2002ApJ...575..996L} have correctly predicted the spin of three BHs in X-ray binaries (XBs). These were measured and reported in \\citet{2006ApJ...636L.113S} and \\citet{2011MNRAS.tmp.1036S}. Using this model, and the estimated BH spins, \\citet{2007ApJ...671L..41B} and \\citet{2011ApJ...727...29M}, have estimated the energies available from spin for a BH to function as a GRB central engine in the known LMXBs (low-mass XBs). These results show that most local LMXBs have too much energy and thus are likely to blow apart the accretion disk feeding the central engine too quickly. Thus, they are likely candidates to explain subluminous LGRBs with very energetic hypernovae, but not Cosmological GRBs. Instead, LMC X-3, may be a more likely candidate for a Cosmological GRB as argued in \\citet{2008ApJ...685.1063B}. There exist three BHs in HMXBs with very large spins \\citep{2008ApJ...679L..37L,2009ApJ...701.1076G,2011ApJ...742...85G}. If these massive BHs ($10 - 20 \\msun$) were born with these spins the energy stored in them is several times $10^{54}$ to $10^{55}$ erg. These formation of these BHs would have likely disrupted the binary had they had such energy available at that time. It is highly likely the spins of these BH were acquired after GCC, however this assumption also carries difficulties  \\citep{2008ApJ...689L...9M,2011MNRAS.413..183M}.  \\citet{2011arXiv1105.4193W} and \\citet{2012ApJ...752...32W} suggest that magnetic torques will spin down the core of the massive star as it differentially rotates after it shrinks to start a new burning stage. This would prevent the star from becoming a collapsar and, thus, from being a GRB progenitor. ",
        "conclusions": "\\label{sec:Concl}  It has been shown that a massive star in a binary that undergoes Case-C mass transfer, a common envelope and tidal synchronization may retain most of the core angular momentum gained through tidal interactions if the magnetic field threading the stellar core and the He envelope is around 100 G.  Furthermore, a magnetar-like field, 10$^{15}$ G, may be easily produced during core collapse (e.g., via a SASI or convective dynamos) given the large amount of available energy during this phase of stellar evolution ($\\sim10^{53}$ to $10^{54}$ erg, as opposed to $10^{47}$ to $10^{50}$ erg to generate the field). This scenario allows a BZ engine to power GRBs and is capable of explaining Cosmological as well as subluminous GRBs.  This model also presents 7 known Galactic XRBs as relics of subluminous GRBs as well as LMC X-3 as a relic of a Cosmological GRB."
    },
    "1402/1402.3323_arXiv.txt": {
        "abstract": " ",
        "introduction": "Existing $N$-body simulations of stellar clusters usually begin with all stars on the zero-age main sequence (ZAMS).  Our motivation to include pre-mainsequence (preMS) evolution in such models goes beyond the desire for completeness. Given the upper mass limit ($M>10\\,M_\\odot$) attributed to conventional star formation \\citep{2000prpl.conf..327S} and because massive stars are mainly found in the denser central regions of star clusters \\citep[e.g.][]{1998ApJ...492..540H}, \\citet{1998MNRAS.298...93B} first suggested collisions between preMS stars as a means to create massive stars.  PreMS stars are larger than their main-sequence counterparts and, although the phase is short-lived, there is a possibility that the increased collision likelihood has an effect.  The journey from molecular gas cloud to protostar to the \\rm{ZAMS} is complicated.  If random motions owing to MHD turbulence are sustained inside a molecular cloud, density inhomogeneities form and quickly collapse under their own gravity to become protostars \\citep*{2005fost}. Typically, when modelling preMS evolution, this self-gravitating fragment of a protostellar cloud is where the models begin \\citep*{tout1999} and is where we define the zero age of the pre-main sequence.  In the 1950s, Henyey constructed the first detailed numerical models of young stars, assuming that preMS stars are radiatively stable \\citep*{1955PASP...67..154H}. This led to nearly horizontal tracks in the Hertzsprung--Russell (H--R) diagram because the stars contract at almost constant luminosity.  \\citet{1961PASJ...13..450H} realised that the $\\rm{H}^-$ opacity (and Kramer's opacity, when the temperature increases) forces the young stars to radiate at nearly constant effective temperature $T_{\\rm{eff}}$ and so to follow nearly vertical tracks in an H--R diagram.  These young stars are in a quasi-hydrostatic equilibrium.  They contract on a time-scale much greater than the free-fall time-scale \\citep*{1961PASJ...13..450H}.  Although protostars are initially homogeneous and isothermal, they collapse non-homogeneously, creating high central density and temperature. It is this that eventually gives rise to the suitable conditions for hydrogen fusion in stars with $M>0.08\\,M_{\\odot}$. Below this mass, the central conditions are never sufficient for hydrogen fusion and the stars collapse to become brown dwarfs \\citep{1963ApJ...137.1121K,1963PThPh..30..460H}.  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=84mm]{fig1.eps} \\caption{An H--R diagram showing the ZAMS in bold and the preMS tracks for a range of masses. The distinction between the vertical Hayashi track and horizontal Henyey track is apparent here. Data is taken from \\citet*{Tout96} and \\citet*{tout1999}.}  \\label{fig:hrdiag} \\end{center} \\end{figure}  Stars with $M \\le 0.7\\,M_{\\odot}$ behave as though they are fully convective and reach the ZAMS at the end of their Hayashi track. When $M \\ge 0.7\\,M_{\\odot}$, the star develops a radiative core which is large enough that it evolves more like Henyey's initial models, turning on to the near-horizontal Henyey track where contraction continues at approximately constant luminosity. This behaviour can be seen in Figure~\\ref{fig:hrdiag}, which shows the calculated preMS tracks modelled by \\citet{tout1999}.  However \\citet{2000prpl.conf..327S} argue that stars more massive than $10\\,M_{\\odot}$ have their contraction disrupted by both strong winds and radiation pressure and thus the conventional theory of protostellar infall fails in this case.  \\citet{2007ARA&A..45..481Z} discussed three competing mechanisms for massive star formation.  These are the turbulent core model of \\citet{2003ApJ...585..850M}, the competitive accretion process of \\citet{1997MNRAS.285..201B} and stellar collisions and merging of stars during their preMS evolution or later.  It is this third case that we investigate here.  An alternative model was investigated by \\citet{baumgardt2011} and it is interesting to compare our results with theirs.  They base their pre-mainsequence evolution on models of accreting stars constructed by \\citet{bernasconi1996}.  The physics used by \\citet{bernasconi1996} is not very different from that used by \\citet{tout1999} and indeed the evolutionary tracks in the H--R diagram are very similar once accretion has ceased and the stars contract down Hayashi tracks.  The major difference between the models we present here and those of \\citet{baumgardt2011} is that all their stars begin as protostellar cores of $0.1\\,\\rm M_\\odot$ and accrete at a constant rate until they arrive at a suitable mass to populate an IMF up to $15\\,\\rm M_\\odot$.  Their stars then undergo a preMS phase in which the logarithm of their radii shrinks linearly with time until they reach the ZAMS, whereupon they begin main-sequence evolution.  Instead we start with a set of pre-mainsequence stars with masses populating a similar IMF up to $8\\,\\rm M_\\odot$ that begin their evolution already fully grown on Hayashi tracks.  Our fitting formulae are a little more complicated because we attempt to to fit both the Hayashi and Henyey phases of pre-mainsequence evolution together.  Because of this our stars remain larger for longer (see section~\\ref{conc}) after any accretion has ceased.  We include this behaviour in cluster models with a variety of initial conditions to test its effect. ",
        "conclusions": "\\label{conc}  We have not attempted to create a completely realistic model of a cluster. The background gas that would exist in a protostellar cloud has been neglected, as have the effects of accretion of this gas \\citep[see the treatment by][]{baumgardt2011}.  Nor have we modelled the effects of stellar accretion discs which would enhance the likelihood of collisions and also change their nature.  Our aim here has been to show that preMS evolution can have important consequences and to identify when the preMS phase should be considered in more detail in the future.  \\begin{figure} \\begin{centering} \\includegraphics[width=84mm]{fig4.eps} \\end{centering} \\caption{The integrated cross-section $I=\\int_{\\tau_{\\rm preMS}} R\\,\\mathrm{d}\\tau$, taken over the preMS lifetime given by equation~(\\ref{eq:tzams}).  The models here are those of \\citet{tout1999}.  The fits of \\citet{baumgardt2011} are integrated from the time their accretion phase stops to when their models reached the ZAMS.  Our cross-sections are relatively larger for the longer lived lower-mass stars but smaller between about 5.5~and $7.5\\,\\rm M_\\odot$.} \\label{fig:xsection} \\end{figure}  It is not straightforward to make a direct comparison with \\citet{baumgardt2011} because they let their stars form by accretion on to $0.1\\,M_{\\odot}$ cores, at constant accretion rates to populate the IMF, whereas we start with an IMF and evolve our protostars as if they were coeval.  However, examining the time-integrated cross-section of our analytic fits compared to theirs (see Figure~\\ref{fig:xsection}) shows that their models systematically underestimate the radii of the preMS stars below $6\\,M_\\odot$ compared to ours.  This is important when looking at collision frequencies. Calculating the time and mass-integrated cross-section \\begin{equation} I_{\\rm{IMF}}=\\int _{0.1\\,M_\\odot}^{4.0\\,M_\\odot} \\int_{\\tau_{\\rm preMS}} \\eta(M) R(M,\\tau)\\,\\mathrm{d}\\tau\\,\\mathrm{d}M, \\end{equation} where $\\eta(M)$ is the Kroupa IMF normalised so that \\begin{equation} \\int _{0.1\\,M_\\odot}^{4.0\\,M_\\odot}\\eta(M)\\,\\mathrm{d}M = 1, \\end{equation} we find that for the original models $I_{\\rm{IMF}}=11.5$, for our analytic fits $I_{\\rm{IMF}}=12.9$ and for the fits of \\citet{baumgardt2011}, $I_{\\rm{IMF}}=7.9$.  Though we should therefore expect to find more collisions in the preMS regime than they do the differences in the models make it almost impossible to identify whether this is actually the case.  The inclusion of gas is a natural next step in these models.  We would expect an interplay between gas expulsion, which decreases the star formation efficiency, accretion rates and the effect of gas accretion on to stars.  Gas accretion not only leads to larger stars and hence larger collision cross-sections but, in clusters, forces contraction of the whole system and thence leads to more merging stars \\citep{1998MNRAS.298...93B,2002MNRAS.336..659B}, an increase in the number of binary stars and early mass segregation \\citep{2011MNRAS.410.2799M}.  We would also expect our preMS stars to form discs while accreting or indeed as a result of tidal disruption during a collision \\citep{2006MNRAS.370.2038D}.  The effect of such protoplanetary discs on the number of collisions, their outcome and the time-scale on which they take place should also be modelled in some way and included in a more detailed $N$-body model.  Star--disc interactions can lead to the disc being stripped away \\citep{disc,2001MNRAS.325..449S} or the formation of binary stars \\citep{1993MNRAS.261..190C} and, because many massive stars live in tight binary systems \\citep{2002ASPC..267..209Z}, it will be important to know whether the preMS phase contributes to this.  In conclusion, although there is much physics that we have not included, our models show that preMS evolution increases both the number of collisions, when the density is sufficiently large, and the amount of mass gained by the final collision remnants.  Most collisions are part of a runaway process so preMS collisions alone are probably insufficient to populate the upper IMF. We identify an initial half-mass cluster density of $10^4\\,M_{\\odot}\\rm{pc}^{-3}$ below which preMS evolution can be neglected because collisions are sufficiently rare.  However caution should be exercised owing to large uncertainty in the early length scales of clusters."
    },
    "1402/1402.6676_arXiv.txt": {
        "abstract": "Effective values for the key helium burning reaction rates, triple-$\\alpha$ and $^{12}$C($\\alpha,\\gamma$)$^{16}$O, are obtained by adjusting their strengths so as to obtain the best match with the solar abundance pattern of isotopes uniquely or predominately made in core collapse supernovae. These effective rates are then used to determine the production of the neutrino isotopes.  The use of effective rates considerably reduces the uncertainties in the production factors arising from uncertainties in the helium burning rates, and improves our ability to use the production of $^{11}$B to constrain the neutrino emission from supernovae. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5675_arXiv.txt": {
        "abstract": "It has been proposed that the first, intermediate-mass ($\\approx 10^{5-6}~M_\\odot$) black holes might form through direct collapse of unpolluted gas in atomic-cooling halos exposed to a strong Lyman-Werner (LW) or near-infrared (NIR) radiation. As these systems are expected to be Compton-thick, photons above 13.6 eV are largely absorbed and re-processed into lower energy bands. It follows that direct collapse black holes (DCBHs) are very bright in the LW/NIR bands, typically outshining small high-redshift galaxies by more than 10 times. Once the first DCBHs form, they then trigger a runaway process of further DCBH formation, producing a sudden rise in their cosmic mass density. The universe enters the ``DCBH era\" at $z \\approx 20$ when a large fraction of atomic-cooling halos are experiencing DCBH formation. By combining the clustering properties of the radiation sources with Monte Carlo simulations we show that in this scenario the DCBH mass density rises from $\\sim 5$~$M_\\odot$ Mpc$^{-3}$ at $z\\sim 30$ to the peak value $\\sim5\\times10^5 M_\\odot$ Mpc$^{-3}$ at $z \\sim 14$ in our fiducial model. However, the abundance of \\textit{active} (accreting) DCBHs drops after $z \\sim 14$, as gas in the potential formation sites (unpolluted halos with virial temperature slightly above $10^4$~K) is photoevaporated. This effect almost completely suppresses DCBH formation after $z\\sim 13$. The DCBH formation era lasts only $\\approx 150$ Myr, but it might crucially provide the seeds of the supermassive black holes (SMBHs) powering $z\\sim6$ quasars. ",
        "introduction": "The formation process of the intermediate-mass black holes (BHs) through direct collapse has been investigated by numerical and theoretical studies \\citep{2002ApJ...569..558O,2003ApJ...596...34B, 2004MNRAS.354..292K, 2006MNRAS.371.1813L,2006MNRAS.370..289B, 2008MNRAS.387.1649B, 2009MNRAS.396..343R,2010MNRAS.402..673B, 2012ApJ...750...66J,2012arXiv1211.0548J,2013MNRAS.433.1607L,2013MNRAS.436.2989L}. Such BHs, due to their much larger mass than the BHs form after the death of Pop III stars (typically several tens $M_\\odot$, see \\citealt{2011Sci...334.1250H}), are good candidates of the seeds of the supermassive black holes (SMBHs) which were observed at $z\\sim6$ \\citep{2001AJ....122.2833F}. In this direct collapse black hole (DCBH) scenario, the formation of a DCBH requires a host halo: (a) with virial temperature above $10^4$~K; (b) exposed to strong Lyman-Werner (LW, photons with energy between 11.2 eV and 13.6 eV) radiation or near-infrared (NIR) radiation with energy above $0.76$~eV \\citep{2007ApJ...665L..85C,2010MNRAS.402.1249S}. Under these conditions, the $\\rm H_2$ can no longer form and the gas in this halo can only cool through Ly$\\alpha$ emission, leading to a quasi-isothermal contraction without fragmentation and finally allowing  a BH seed with mass $\\sim10^4 - 10^6~M_\\odot $ to form at the center of the halo.  If such DCBHs largely formed at $z\\simgt 13$ and are all Compton-thick, with a typical mass $\\sim10^6~M_\\odot$ and typical accretion time scale $\\sim 30~$Myr, they might be excellent candidate sources of the fluctuations of the source-subtracted near infrared background (NIRB) \\citep{2005Natur.438...45K,2007ApJ...654L...5K, 2007ApJ...657..669T,2011ApJ...742..124M,2012ApJ...753...63K,2012Natur.490..514C} which is hard attributed to normal galaxies \\citep{2012ApJ...752..113H,2012ApJ...756...92C, 2013MNRAS.431..383Y}. As proposed in \\citet{2013MNRAS.433.1556Y}, DCBHs also naturally generate a cosmic X-ray background (CXB) that is correlated with the NIRB, which is observed by \\citet{2013ApJ...769...68C}.  The minimum LW intensity necessary to sustain the formation of the DCBH has been somewhat debated in recent literatures. Early studies favored values as high as $\\sim10^3$ (in units of $10^{-21}$erg s$^{-1}$cm$^{-2}$Hz$^{-1}$sr$^{-1}$) for a $10^4$~K blackbody spectrum, which represents a typical spectrum of Pop II stars; and $\\sim10^{4-5}$ for a $10^5$~K blackbody spectrum, which characterizes metal-free stars \\citep{2000MNRAS.317..175S,2001ApJ...546..635O,2002ApJ...569..558O,2003ApJ...596...34B}. The critical intensity for the $10^4$~K blackbody spectrum is lower because in this case, the main mechanism that quenches the $\\rm H_2$ formation is not to directly dissociate the $\\rm H_2$ themselves, but to dissociate the $\\rm H^{-}$ \\citep{2007ApJ...665L..85C,2012MNRAS.425L..51W} \\--- the necessary intermediary for the $\\rm H_2$ formation \\--- through the NIR radiation above $0.76$~eV. However, a more detailed work by \\citet{2010MNRAS.402.1249S} revised the value downward: e.g. for a $10^4$~K blackbody spectrum, it is $\\sim 30-300$. The critical LW intensity could be even lower if a more refined treatment of self-shielding factor is implemented (\\citealt{2011MNRAS.418..838W}, Z. Haiman, personal communication).  Usually the probability for a halo to be embedded in a radiation field with super-critical LW intensity\\footnote{Throughout this paper, when we are talking about the ``super-critical LW intensity\", we refer to either the LW radiation above a critical intensity \\--- in case the $\\rm H_2$ cooling is suppressed due to the direct dissociation of $\\rm H_2$, or super-critical NIR intensity which can be characterized by the intensity at LW band when the spectrum is already known \\--- in case the $\\rm H^-$ dissociation plays the main role. } is very small if normal galaxies and/or Pop III stars are the only sources of LW photons \\citep{2008MNRAS.391.1961D,2012MNRAS.425.2854A,2013MNRAS.432.3438A}; in this case  DCBHs are very rare. However, \\citet{2013MNRAS.433.1556Y} pointed out that a Compton-thick BH could be very luminous below 13.6 eV due to the  reprocessed nebular emission. This means that, once the first DCBHs form, they could ignite a runaway process  favoring the formation of more DCBHs and leading to a steep rise in the DCBH abundance; we dub this well-defined period during cosmic evolution as the ``DCBH era'', lasting for about 150 Myr.  The probability that a halo sees a LW intensity above the critical value can be  calculated by Monte Carlo simulations \\citep{2008MNRAS.391.1961D}, in which the LW radiation from previously formed DCBHs is explicitly taken into account. With this probability, and the probability of a halo being unpolluted, we get the formation rate of the DCBH and build the picture of the rise process of its abundance. The process comes to an end when the host atomic-cooling halos ($T_{\\rm vir} \\simgt 10^4$~K) are rapidly photoevaporated by ionizing photons. Photoevaporation drastically reduces the DCBH formation rate and almost completely quenches their formation for  $z\\simlt 13$. This is the basic idea of this work. We aim at building a solid picture of the rapid rise (and fall) of the DCBH formation rate during the DCBH era.  The layout of this paper\\footnote{Throughout the paper, we use $M_{\\rm T3}$, $M_{\\rm T4}$ and $M_{\\rm 2T4}$ to denote the mass of halos with $T_{\\rm vir} = 10^3$~K, $10^4$~K and $2\\times10^4$~K respectively at the corresponding redshift. The cosmological parameters are the same as used in \\cite{2011MNRAS.414..847S}:  $\\Omega_m$=0.26, $\\Omega_\\Lambda$=0.74, $h$=0.73, $\\Omega_b$=0.041, $n=1$ and $\\sigma_8$=0.8. We use the the transfer function in \\cite{1998ApJ...496..605E}.} is as follows: in Sec. \\ref{methods} we introduce the clustering of the spatial distribution of halos at high redshifts, the LW luminosity of these sources and the various mechanisms that influence the formation of DCBHs. We give results in Sec. \\ref{results}, the conclusions are presented in Sec. \\ref{conclusions}. A list of the mostly used symbols is given in Tab. \\ref{symbols}. The calculation of the probability that a halo sees a super-critical LW intensity is introduced in Appendix \\ref{AppA}.  \\begin{table} \\begin{center} \\caption{ {List of the most used symbols.} } \\begin{tabular}{ll} \\hline \\hline Symbol & Definition \\\\ \\hline {\\small $ L_{\\rm LW}^{\\rm gal}(M)$} &\\multirow{1}{13cm}{\\small LW luminosity of a normal galaxy hosted by a halo of mass $M$} \\\\ {\\small $L_{\\rm LW}^{\\rm BH}(M_{\\rm BH})$}&\\multirow{1}{13cm}{\\small LW luminosity of a DCBH with mass $M_{\\rm BH}$}\\\\ {\\small $l_{\\rm LW}^{\\rm pop3}$} &\\multirow{1}{13cm}{\\small LW luminosity (per stellar mass) of Pop III stars }\\\\ {\\small $J_{\\rm LW}^{\\rm bg}$}&\\multirow{1}{13cm}{\\small Background LW intensity} \\\\ {\\small $J_{\\rm LW}^{\\rm pop3}$}&\\multirow{1}{13cm}{\\small Background LW intensity from Pop III stars}\\\\ {\\small $J_{\\rm LW}^{\\rm gal}$}&\\multirow{1}{13cm}{\\small Background LW intensity from normal galaxies }\\\\ {\\small $J_{\\rm LW}^{\\rm BH}$} & \\multirow{1}{13cm}{\\small Background LW intensity from DCBHs }\\\\ {\\small $J_{\\rm LW}^{\\rm crit}$}&\\multirow{1}{13cm}{\\small Critical LW intensity for DCBH formation} \\\\ {\\small $f_\\star^{\\rm gal}$} &\\multirow{1}{13cm}{\\small Star formation efficiency of normal galaxies}\\\\ {\\small $f_\\star^{\\rm pop3}$} &\\multirow{1}{13cm}{\\small Star formation efficiency of Pop III stars} \\\\ {\\small $M_{\\rm crit,0}^{\\rm pop3}$}&\\multirow{1}{13cm}{\\small Critical minihalo mass for Pop III stars formation without LW feedback}\\\\ {\\small $M_{\\rm crit}^{\\rm pop3}$}&\\multirow{1}{13cm}{\\small Critical minihalo mass for Pop III stars formation with LW feedback}\\\\ {\\small $M_{\\rm min}^{\\rm ion}$}&\\multirow{1}{13cm}{\\small Critical halo mass for galaxy/DCBH formation under photoevaporation feedback} \\\\ {\\small $\\rho_{\\rm pop3}$}&\\multirow{1}{11cm}{\\small Mass density of Pop III stars }\\\\ {\\small $\\rho_{\\rm BH}$}&\\multirow{1}{11cm}{\\small Mass density of active DCBHs}\\\\ {\\small $\\rho_{\\rm BH}^{\\rm cum}$}&\\multirow{1}{11cm}{\\small Cumulative DCBH mass density}\\\\ {\\small $p_{\\rm g}$}&\\multirow{1}{11cm}{\\small Probability of genetic enrichment} \\\\ {\\small $p_{\\rm J}$}&\\multirow{1}{11cm}{\\small Probability that a halo sees a super-critical LW intensity}\\\\ {\\small $p_{\\rm W}$}&\\multirow{1}{13cm}{\\small Probability that a halo being enriched by metals carried by galactic winds} \\\\ {\\small $p_{\\rm J-W}$}&\\multirow{2}{13cm}{\\small Probability that a halo sees a super-critical LW intensity but without being polluted by galactic winds} \\\\ \\smallskip\\\\ {\\small $P(>z,M|M_0,z_0)$}&\\multirow{2}{13cm}{\\small For a halo with mass $M_0$ at $z_0$, the probability that its most massive progenitor with mass $M$ formed before $z$} \\smallskip \\\\ \\smallskip \\\\ {\\small $\\frac{dP}{dz}(z,M|M_0,z_0)$}&\\multirow{1}{13cm} {\\small Redshift derivative of above probability}\\\\ \\smallskip \\\\ \\hline \\hline \\end{tabular} \\label{symbols} \\end{center} \\end{table} ",
        "conclusions": "We have studied the formation and evolution of DCBHs at high ($z$ \\simgt 13) redshifts. The most striking result is that we found that a well-defined DCBH formation era might occur in the early universe extending approximately for 150 Myr from $z\\sim 20$ to $z\\sim 13$. The intermediate masses ($M\\sim10^{6} M_\\odot$) of DCBHs make them excellent seeds for the subsequent rapid growth of the SMBH observed at $z$ \\simgt 6. In addition, such BHs could be very abundant in this era and might leave a clear signature in the observed fluctuations of the source-subtracted NIRB \\citep{2013MNRAS.433.1556Y}.  During such era the abundance of active DCBH grows, reaches a peak and eventually comes to a sudden halt when the required conditions for their formation disappear. These conditions include the availability of pristine gas in atomic-cooling halo ($T_{\\rm vir} \\simgt 10^4$~K) embedded in a sufficiently strong LW or NIR radiation field. Whether and for how long these conditions can be met during cosmic evolution depends on a complex network of negative and positive feedbacks that we have described in detail.  A key result is that the (so far neglected) positive feedback of DCBH on their own formation fosters a rapid multiplication of the number of DCBHs. This is possible because DCBHs are likely Compton-thick systems (see \\citealt{2013MNRAS.433.1556Y}), they are very bright in the rest-frame UV/NIR band, typically outshining small high-redshift galaxies by more than 10 times; in their surroundings the LW flux is therefore largely enhanced with respect to the level provided by galaxies, thus enabling additional DCBHs to form. DCBH formation spreads very rapidly, similarly to a flame in a haystack. Their mass density raises from $\\sim 5$~$M_\\odot$Mpc$^{-3}$ at $z \\sim 30$ to the peak value $\\sim5\\times10^5$~$M_\\odot$Mpc$^{-3}$ (corresponding to a density parameter $\\Omega_{\\rm BH} \\sim 3\\times10^{-6}$)  at $z \\sim 14$ in our fiducial model.   The DCBH formation era however does not last very long. This is due to a combination of several negative feedback effects that we have discussed in Sec. \\ref{methods}. In addition to the effect of a decreasing bias of DCBH host halos with time, metal pollution, either due to heavy elements inherited by the progenitor halos or carried by supernova winds coming from neighbors, is perhaps an obvious way to quench DCBH formation. However, we find that collectively these effects are not able to quench the DCBH widespread formation before $z\\sim6 - 7$. Instead, the rapid DCBH birthrate drop is caused by photoevaporation of the gas from the gravitational well of the smallest atomic-cooling halos.  This effect reduces the formation rate of the DCBH, almost completely suppressing their formation for  $z\\simlt 13$. This is the reason why DCBH formation era lasts for a relatively short time interval during cosmic evolution. However, as they are Compton-thick, their ionizing photons are fully reprocessed into lower energy radiation before escaping into the IGM and, as such, they play no active role in driving reionization.  The above scenario is robust against variations of the free parameters of the model, for which we have explored (see Sec. \\ref{modvar}), including the critical LW intensity threshold for DCBH formation, the strength of the genetic enrichment, and the influence of the assumed minimum distance between two halos. In spite of these, we do not observe substantial deviations from the evolutionary trend described by the fiducial case, witnessing a DCBH formation history firmly controlled by feedback(s) action.  Our work is based on the most popular assumption that DCBH can only form under specific conditions, i.e., the gas is metal-free and the $\\rm H_2$ formation is suppressed by strong external radiation. It is possible that even in the presence of metals or $\\rm H_2$, if the accretion timescale is smaller than the fragmentation timescale, DCBH may still form  \\citep{2009ApJ...702L...5B, 2010Natur.466.1082M,2014MNRAS.437.1576B}; particularly because the supersonic turbulence suppresses the fragmentation \\citep{2005ApJ...630..250K,2007ApJ...656..959K}. Albeit interesting, this scenario still needs to overcome a number of open questions that have been discussed in a recently published paper by two of the present authors \\citep{2013MNRAS.434.2600F}. We cannot exclude different DCBH formation channels with respect to the canonical one explored here, however, the photoevaporation of haloes with circular velocity \\simlt 30~km/s, which is discussed in this paper, might in practice make these channels only have modest influence on our conclusions. We also note that purely forming DCBH in metal free, UV illuminated halos already produces a BH cosmic mass density comparable to the present-day one.  We finally comment on a possibly intriguing implication of our model. The end of the DCBH era leaves behind a large number of ``fossil'' DCBH seeds ($\\Omega_{\\rm BH} \\sim 10^{-5}$). Therefore to have sufficient number of seeds for SMBHs is not a problem any more in this scenario (see also \\citealt{2012MNRAS.425.2854A}). Some of DCBHs are ready to be embedded in other galaxies and in SMBHs. While only a small fraction of these BHs will continue to grow following merging events and possibly becoming SMBHs, the majority of them will evolve passively. Hence a large number of intermediate-mass black holes could be present in local universe, which has escaped so far detection. In the future it will be  interesting to consider suitable strategies to find them. It is also tempting to speculate that DCBH gas--devoid, low-mass dark matter halo hosts might be related with the large missing halo population predicted by $\\Lambda$CDM models at $z=0$."
    },
    "1402/1402.5219_arXiv.txt": {
        "abstract": "We present recent results from our searches of 21-cm absorption using GBT, GMRT, VLBA and WSRT to trace the evolution of cold gas in galaxies. Using $\\sim$130 sight lines with 21-cm absorption measurements, we find that within the measurement uncertainty, the 21-cm detection rate in strong Mg~{\\sc ii} systems is constant over $0.5<z<1.5$. Since stellar feedback processes are expected to diminish the filling factor of CNM over $0.5< z <1$, this lack of evolution in the 21-cm detection rate in Mg~{\\sc ii} absorbers is intriguing. Further, we find that if majority of 21-cm absorbers arise from DLAs then the cross-section of 21-cm absorbing gas i.e. cold neutral medium amongst DLAs has increased from $z=3.5$ to $z=0.5$. In a sample of 13 $z>2$ DLAs with both 21-cm and H$_2$ (another tracer of cold gas) absorption measurements, we report two new H$_2$ detections and find that in 8/13 cases neither 21-cm nor H$_2$ is detected. This confirms that the H~{\\sc i} gas in $z>2$ DLAs is predominantly warm.  Interestingly, there are two cases where 21-cm absorption is not detected despite the presence of H$_2$ with evidence for the presence of cold gas.  This can be explained if H$_2$ components seen in DLA are compact ($\\le$15\\,pc) and contain $\\le$10\\% of the total N(H~{\\sc i}). We briefly discuss results from our ongoing survey to identify 21-cm absorbers at low-$z$ to establish connection between 21-cm absorbers and galaxies, and constrain the extent of absorbing gas. ",
        "introduction": "\\label{s:intro}  It is well known that physical conditions in the diffuse interstellar medium (ISM) of galaxies are influenced by various radiative and mechanical feedback processes associated with in-situ star formation. Therefore, volume-filling factors of different phases of gas in a galaxy are expected to depend on its star formation history. Of particular interest is the evolution of the volume-filling factor of cold neutral medium (CNM) phase that also serves as a gaseous reservoir for star formation in galaxies. Systematic searches of high-$z$ intervening 21-cm absorbers in samples of Mg~{\\sc ii} systems and damped Ly$\\alpha$ systems (DLAs) towards QSOs to measure CNM filling factor of galaxies have resulted in detections of 21-cm absorption towards $\\sim$10-20\\% of Mg~{\\sc ii} systems and DLAs. However, establishing a connection between the redshift evolution of 21-cm absorbers and global star formation rate (SFR) density is not straight forward due to (i) small number statistics of 21-cm absorbers; (ii) ambiguities regarding the origin of absorbing gas; and (iii) issues related to the small scale structure in absorbing gas and the extent of radio source. Here we present a summary of our recent results to constrain the evolution of cold gas in galaxies using 21-cm absorption and address issues related to partial coverage of absorbing gas. ",
        "conclusions": ""
    },
    "1402/1402.2325_arXiv.txt": {
        "abstract": "Galaxy clusters contain much more metal per star, typically 3 times as much, than is produced in normal galaxies. We set out to determine what changes are needed to the stellar mass function and supernovae rates to account for this excess metal. In particular, we vary the Type Ia supernovae rate, IMF slope, upper and lower mass cutoffs, and the merger rate of massive stars. We then use existing simulation results for metal production from AGB stars, Type Ia SNe and core-collapse SNe to calculate the total amount of each element produced per solar mass of star formation. For models with very massive stars, we also include metal production from pair-instability supernovae (PISNe). We find that including PISNe makes it much easier to increase the amount of metal produced per stellar mass. Therefore a separate population of high-mass stars is not needed to produce the high amounts of metal found in galaxy clusters. We also find that including at least some PISNe increases the abundance of intermediate-mass elements relative to both oxygen and iron, consistent with observations of ICM abundances. ",
        "introduction": "\\label{sec:introduction}  The hot intra-cluster medium (ICM) of galaxy clusters is not made of pristine primordial gas, but instead is enriched in iron to $30\\%$ solar or more \\citep{edge91,yamashita92,mushotzky97}, including super-solar iron abundances in the Centaurus cluster \\citep{sanders06,lovisari11}. Abundances of other elements are also enhanced. The ICM in a galaxy cluster typically contains 85\\% to 95\\% of the total baryons in the cluster, with stars accounting for just 5\\% to 15\\% \\citep{gonzalez07,andreon10}. The cluster galaxies themselves are slightly enriched relative to field galaxies of the same types \\citep[e.g.][]{ellison09}.  Overall, the total amount of metal per star in a galaxy cluster is around 2-6 times more than is found in the Milky way. Put another way, if all the metal in galaxy clusters were contained in the cluster galaxies, these galaxies would have metallicities 2-6 times solar.  This is very difficult to explain with normal star formation and stellar evolution. Typically, a galaxy or any collection of stars and gas should saturate at an average metallicity close to solar \\citep{tinsley80,pagel97}. Although some stars will have super-solar metallicity, this is because they formed from gas enriched by previous generations of lower metallicity stars, so the average metallicity is still solar. For 1 M$_{\\sun}$ of stars formed, about $0.02$ M$_{\\sun}$ of metal will be produced via nucleosynthesis in AGB stars and supernovae.  Stars not in galaxies, which produce the intra-cluster light, are not sufficient to account for the excess amount of metal if they have a normal IMF, as they make up only about 5\\% to 20\\% of the stars in clusters \\citep{krick07}.  Numerous previous studies \\citep[e.g.][]{loewenstein01,portinari04,loewenstein06,loewenstein13} have rejected the notion that the metal in galaxy clusters can come from a standard initial-mass function (IMF). They concluded that either a separate population of high-mass stars \\citep{loewenstein01,loewenstein06} or a top-heavy IMF \\citep{portinari04,fabjan08} is necessary to produce enough iron and other metals through additional Type Ia and core-collapse SN. \\citet{bregman10} concluded that the majority of metals in a cluster must form outside of cluster galaxies, in population III stars.  There is observational evidence for an increased rate of Type Ia supernovae in galaxy clusters. \\citet{mannucci08} and \\citet{sharon07} found that Type Ia rates in cluster elliptical galaxies are about 3 times the rate in field ellipticals. \\citep[However, see][]{sand12,maoz12}. Assuming the increased Type Ia rate reflects a true increase in the number of Type Ia SNe per solar mass of stars formed, it would be insufficient to account for the increased iron amounts in clusters \\citep[see sect.~\\ref{sec:typeIa},][]{loewenstein06}. \\citet{mannucci08} found the the core-collapse SNe rate in cluster and field galaxies is consistent with being the same, indicating there cannot be a dramatically different IMF in cluster galaxies.  Evidence from low-mass stars also indicates the IMF in galaxy clusters cannot be dramatically different. There appears to be an excess number of low-mass stars in large elliptical galaxies \\citep[e.g.][]{treu10,auger10,vandokkum10}, possibly indicating a time-variable IMF \\citep{weidner13a}. If there is a large excess population of high-mass stars in clusters, it would have to be above 1 M$_{\\sun}$ only, and not just a flattening of the low-mass IMF slope.   Previous studies do not include the impact of vary massive stars ($>140$ M$_{\\sun}$) that produce pair-instability supernovae (PISNe). The IMF is typically truncated at a fixed upper mass of 100 to 150 M$_{\\sun}$ \\citep{portinari04,loewenstein06,loewenstein13}, approximately the mass of the largest stars in the Milky Way. Alternately, larger stars could be included but assumed to lose a substantial fraction of their mass and produce core-collapse supernovae. In a typical IMF, the number of very massive stars is so small the a few extra core-collapse SNe will not substantially alter the total metal production. However, if very massive stars are present and they each produce a PISN, they would produce substantially more metal. A PISN can, for example, produce a few 10's of solar masses of iron \\citep{heger02,galyam09} rather than a few tenths of a solar mass of iron in a core-collapse SN \\citep[e.g.][]{heger10}.     \\subsection{Evidence star formation is different in clusters}  There are several reasons to think very high-mass stars may be able to form in galaxy clusters. Cluster galaxies, particularly cD galaxies, have very large star forming regions, up to several million solar masses \\citep[see, e.g.][]{canning10}. They also have very high star formation rates of 10 to 100 M$_{\\sun}$ per year in nearby clusters \\citep[e.g.][]{mcnamara89,odea10} and up to 740 M$_{\\sun}$ per year at high redshift \\citep{mcdonald12}. Very high-mass stars may only form in very large star forming regions. This is particularly true if stellar masses in a star forming region are optimally sampling the IMF, rather than randomly sampling it \\citep[see][]{weidner13}. In the Milky Way, the largest star forming regions do contain the highest-mass stars \\citep{weidner13}. In the local group, the largest stars and the largest number of O stars are found in R136, part of the 30 Doradous complex in the LMC, the largest nearby star forming region \\citep{massey98}.  The largest star in R136 is $265$ M$_{\\sun}$ ($320$ M$_{\\sun}$ initial mass), much larger than any star known in the Milky Way \\citep{crowther10}, and well above the mass cutoff assumed in a typical IMF. It is unknown if this star formed as a single star or from the merger of two massive stars \\citep{banerjee12}, but in either case is shows that very high-mass stars can form in very massive star forming regions.   cD galaxies in clusters also have a much larger number of globular clusters than the Milky Way. Our Galaxy has about 150-200 globular clusters, while M87, with about 10 times the stellar mass, has about 12,000 \\citep{tamura06}. The specific frequency of globular clusters in cD galaxies is about 20 times the frequency in the Milky Way and 3-6 times the frequency in elliptical galaxies \\citep{harris91}. Globular clusters likely began their lives as massive star forming regions, so something about star formation in cD galaxies must allow proportionally more of these regions to exist over the life of the Universe. This may indicate the typical star formation region size is always larger in cluster galaxies than in the Milky Way.   Aside from the size of the star forming regions, galaxy clusters, at least in their inner regions, are at very high pressure. In the inner 100 kpc, the typical ICM pressure in a galaxy cluster is 5 to 50 times the ISM pressure in the Milky Way \\citep[e.g][]{arnaud10}. This presumably means an increased pressure in star forming regions as well, which could lead to more rapid accretion onto stellar cores and therefore a higher upper mass limit.     Very large star-forming regions may also have a slightly flatter IMF slope. For example, the Arches star cluster appears to have a slope of $-2.1$ rather than the standard $-2.3$ \\citep{espinoza09}, and NGC 3603 likely has an even flatter slope, measured as $-1.74^{+0.62}_{-0.47}$ by \\citet{harayama08}. Although this is only slightly flatter, the slope has the largest effect as the highest masses, so even a small change can significantly increase the number of very high mass stars.   Even if the {\\it initial} mass function in cluster galaxies has the same mass cutoff as the Milky Way, very high mass stars may still form. Galaxies in clusters, and very massive star forming regions, have very dense stellar environments. This will increase the rate of stellar interactions compared to field galaxies, and a hardening of binary systems. \\citet{sana13} found that 50\\% of the O stars in 30 Doradus will exchange mass with a companion during their lifetime. This could lead to mergers of high mass stars to produce very high mass stars \\citep{banerjee12}. Type Ia supernovae are believed to be produced from binary interactions, either accretion onto a white dwarf from a companion or a merger of two white dwarfs, so an increase in stellar interactions may also explain the higher Type Ia SNe rate seen in cluster galaxies \\citep{mannucci08}.  Pressures and interaction rates will be highest in the central regions of galaxy clusters, and particularly in cD galaxies. However, this accounts for only a small fraction of the total galaxy cluster mass (about 10\\% of stars are in the cD galaxy). If extra metal is produced only in this region, the excess relative to ordinary star formation would have to be very large, and the metal would need to be transported to the outer region of the cluster. This could possibly be accomplished through AGN driven outflows or sloshing of cluster gas. In galaxy clusters without a cD galaxy, or with lower density non-cool cores, it may be less likely to have a significantly modified IMF, and therefore be more difficult to produce excess metals.  If enhanced metal production compared to the Milky Way is a general feature of elliptical galaxies, and not just galaxies in the inner regions of clusters, we might expect evidence from observations of metals in ellipticals and in the IGM, particularly if PISN are responsible. Stellar abundances in the centers of ellipticals, determined through optical observations of integrated starlight, indicate that the total amount of metal is enhanced relative to solar by a factor about 2, and that $\\alpha$ elements (typically Mg) are enhanced relative to iron by about 1.5 \\citep[see, e.g.,][]{trager00, trager00b, tang09}. These enhancements also correlate with galaxy mass. The IGM, though far below solar metallicity, has an apparent enhancement of Si relative to carbon ([Si/C]$=0.77 \\pm 0.05$ rather than $\\approx 0.5$), possibly indicating a significant contribution from PISNe \\citep{aguirre04}.  We aim to explore what combination of stellar mass function and supernovae rates, including PISNe, are needed to produce the high amount of metal per star seen in galaxy clusters. In section~\\ref{sec:base_models} we explain our methods for calculating total metal production and compare our results using standard IMF base models to observed solar metallicity values. In section~\\ref{sec:results} we explore how various modifications to our base models increase metal yield. In section~\\ref{sec:compare_obs} we compare the results of our models to observed abundances of various elements in galaxy clusters. In section~\\ref{sec:kobayashi} we examine how our results change if we use an alternative set of core-collapse supernova models. Finally in section~\\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": "  1. Including contributions from pair-instability supernova significantly increases total metal production. This is particularly true when combined with an increased type Ia SN rate and a slightly flattened IMF slope. The PISNe can be produced either by increasing the upper mass cutoff of the IMF or by mergers between stars. A spatially or temporally distinct population of high-mass stars, e.g. a very top-heavy IMF for intra-cluster stars or a large amount of Population III star formation, is not necessary to produce the high amounts of metal found in galaxy clusters.   2. PISN are necessary to simultaneously produce enhancements of intermediate-mass elements relative to oxygen and iron. Each means of increasing metal production in clusters create a distinct abundance signature. An enhancement of intermediate elements relative to oxygen can be produced by excess type Ia SNe, but with a depletion relative to iron. The opposite will occur if there is an excess of core-collapse SNe (i.e. for a flattened IMF). Models with an excess of both core-collapse and Type Ia SNe, or with low-mass stars removed, will produce approximately solar abundance ratios. Only by including PISN can the amount of intermediate elements be increased relative to both oxygen and iron.  3. Models that simply extend the high-mass slope of the IMF to a cutoff mass above the PISN range over produce intermediate elements relative to oxygen and iron. Models with PISNe from mergers of 10\\% to 30\\% of stars above 65 M$_{\\sun}$ appear to produce better fits. This could indicate that there are fewer stars then expected for a continuous IMF in the right mass range (140 to 260 M$_{\\sun}$), that the upper mass cutoff of the IMF is still below 260 M$_{\\sun}$ in galaxy clusters, or that an initial mass higher than 140 M$_{\\sun}$ is required to produce a PISN. It could also be that not all galaxies in clusters produce PISN. For example, it may be that only massive galaxies or only galaxies near the center of the cluster are able to produce stars large enough for a PISN to occur.  4. We expect 1 PISN to occur for every 60 to 600 core-collapse supernovae in galaxy clusters. For a cluster with a star formation rate of 100 M$_{\\sun}$/yr, this would translate to 1 PISN every 20 to 600 years, depending on the IMF. This assumes the number of PISN per stellar mass formed is constant over the life of a cluster. It is possible PISN could only occur at and before the time of peak star formation in the cluster, when average metallicity would have been lower. This would, however, require more PISNe per stellar mass formed at these times.  The largest uncertainties in our models come from supernova simulations results used as input. None of the simulations of type Ia, core-collapse, or pair-instability supernovae are completely realistic, e.g., none are 3D simulations, none naturally produce a SN, etc. Better models might change the amounts of metals produced by different supernovae, or change the relative amounts of specific elements, but so long at the general trends are the same this should not qualitatively change our results. We also do not consider the metallicity history of the cluster galaxies in our models, aside from using a fixed range of metallicities for AGB stars.  It is also unknown under what conditions PISNe will be created. The simulation results we use \\citep{heger02} are for non-rotating stars with an initial metallicity of zero. Stars with non-zero metallicity may loose too much mass for a PISN to occur, or the initial mass lower-limit for PISN may be higher. Alternatively, more recent work \\citep{chatzopoulos12} has shown that PISN can occur a much lower mass ($\\approx 65$~M$_{\\sun}$ rather than $140$~M$_{\\sun}$) for rapidly-rotating stars. Real PISNe may occur at different masses and produce very different amounts of each element that the values we assume here. In particular, the abundance of intermediate-mass elements is very sensitive to the number of PISNe in our models, and is therefore very dependent on the PISNe models we use. However, so long as the general result that PISNe produce enhanced amounts of intermediate elements relative to oxygen and iron remains true, the conclusion that PISNe are necessary in galaxy clusters will remain true."
    },
    "1402/1402.1711_arXiv.txt": {
        "abstract": "We present a study of the spatially resolved radio continuum--star formation rate (RC--SFR) relation using state-of-the-art star formation tracers in a sample of 17 THINGS galaxies. We use SFR surface density ($\\Sigma_{\\rm SFR}$) maps created by a linear combination of \\emph{GALEX} far-UV (FUV) and \\emph{Spitzer} $24~\\mu\\rm m$ maps. We use RC maps at $\\lambda\\lambda$ 22 and 18~cm from the WSRT SINGS survey and H$\\alpha$ emission maps to correct for thermal RC emission. We compare azimuthally averaged radial profiles of the RC and FUV/mid-IR (MIR) based $\\Sigma_{\\rm SFR}$ maps and study pixel-by-pixel correlations at fixed linear scales of $1.2$ and $\\rm 0.7~kpc$. The ratio of the integrated SFRs from the RC emission to that of the FUV/MIR-based SF tracers is $\\mathscr{R}_{\\rm int} = 0.78\\pm 0.38$, consistent with the relation by Condon. We find a tight correlation between the radial profiles of the radio and FUV/MIR-based $\\Sigma_{\\rm SFR}$ for the entire extent of the disk. The ratio $\\mathscr{R}$ of the azimuthally averaged radio to FUV/MIR-based $\\Sigma_{\\rm SFR}$ agrees with the integrated ratio and has only quasi-random fluctuations with galactocentric radius that are relatively small (25\\%). Pixel-by-pixel plots show a tight correlation in log-log diagrams of radio to FUV/MIR-based $\\Sigma_{\\rm SFR}$, with a typical standard deviation of a factor of two. Averaged over our sample we find $\\rm (\\Sigma_{SFR})_{RC}\\propto (\\Sigma_{SFR})_{hyb}^{0.63\\pm0.25}$, implying that data points with high $\\Sigma_{\\rm SFR}$ are relatively radio dim, whereas the reverse is true for low $\\Sigma_{\\rm SFR}$. We interpret this as a result of spectral aging of cosmic-ray electrons (CREs), which are diffusing away from the star formation sites where they are injected into the interstellar medium. This is supported by our finding that the radio spectral index is a second parameter in pixel-by-pixel plots: those data points dominated by young CREs are relatively radio dim, while those dominated by old CREs are slightly more RC bright than what would be expected from a linear extrapolation. We studied the ratio $\\mathscr{R}$ of radio to FUV/MIR-based integrated SFR as a function of global galaxy parameters and found no clear correlation. This suggests that we can use RC emission as a universal star formation tracer for galaxies with a similar degree of accuracy as other tracers, if we restrict ourselves to global or azimuthally averaged measurements. We can reconcile our finding of an almost linear RC--SFR relation and sub-linear resolved (on 1~kpc scale) RC--$\\Sigma_{\\rm SFR}$ relation by proposing a non-linear magnetic field--SFR relation, $B\\propto \\rm SFR_{hyb}^{0.30\\pm 0.02}$, which holds both globally and locally. ",
        "introduction": "\\label{sec:introduction} Radio continuum (RC) emission in galaxies stems from massive stars, allowing us to use it as a star formation (SF) tracer. Ionizing UV radiation from massive stars creates \\ion{H}{2} regions where thermal electrons give rise to RC radiation through the process of {\\em bremsstrahlung} or thermal (free--free) emission. These same stars, upon reaching the end of their lives, are the source of cosmic-ray electrons (CREs) that are accelerated in supernova shock waves. When these CREs encounter a magnetic field, they generate non-thermal (synchrotron) emission. If all CREs lose their energy exclusively {\\em within} the galaxy, it can be considered a {\\em calorimeter} \\citep{voelk_89a,lisenfeld_96a}. More realistically, a galaxy is a leaky box, and models become far more involved \\citep{bell_03a,lacki_10a}. The relative success of theoretical models lends confidence that massive SF and RC emission are indeed closely tied together.  RC can act as tracers for the star formation rate (SFR) with the added advantage over other tracers---such as the optical Balmer H$\\alpha$ line emission, far-UV (FUV) emission as measured by the \\emph{GALEX} satellite, and hybrid combination of these \\citep[e.g.,][]{leroy_12a}---that they are unaffected by dust attenuation. The advantage offered by probing SF via radio astronomical means is that this does not rely on the availability of (cryogenically cooled) IR satellites. The motivation of this paper is to deepen our understanding of the relation between the SFR and the RC luminosity of normal star-forming galaxies, hereafter the RC--SFR relation, as originally explored by \\citet{condon_92a}.  RC emission observed at 33~GHz is an almost ideal SF tracer \\citep{murphy_11a}. The reason is that thermal free--free emission is the dominating emission mechanism in the 20--40~GHz range. However, galaxies can be a factor of 10 brighter at lower frequencies such as between 1 and 2~GHz than at 33~GHz, where the non-thermal emission dominates. This frequency range is hence potentially more valuable, especially for cosmological studies. For this to work, the non-thermal component needs to be calibrated as successfully as the thermal RC against other SF indicators \\citep[see][for a recent review of SF tracers]{kennicutt_12a}. A properly calibrated RC--SFR for non-thermal continuum emission would offer a powerful method to probe the cosmic SFR, initially with the Very Large Array (VLA) and with Square Kilometre Array (SKA) precursors, and eventually with the SKA \\citep[e.g.,][]{murphy_09a,lacki_10b}. At the highest redshifts inverse Compton losses could dominate to such an extent that only the (weaker) thermal RC is expected to survive, although opinion on this remains divided \\citep[e.g.,][]{murphy_09a,lacki_10b,schleicher_13a}.  A semi-empirical model was introduced by \\citet{condon_92a} to quantify the RC--SFR relation. It assumes that the RC luminosity is proportional to the CRE production rate, which itself is proportional to the supernova rate and hence to the SFR. The Condon relation is a \\emph{calorimetric model}, which requires at a minimum the following two assumptions to hold \\citep{voelk_89a,lisenfeld_96a}: (1) galaxies act as a calorimeter both for the CREs and for the ionizing UV-photons: and (2) although CREs have competing energy losses to synchrotron radiation such as inverse Compton (IC) radiation, ionization, bremsstrahlung, and adiabatic losses, the fraction of these additional losses to the synchrotron radiation should remain constant. The assumption of a CRE calorimeter does hold in starburst galaxies \\citep[e.g.,][]{paglione_12a,yoast-hull_13a}, but probably not in normal spiral galaxies. Therefore, as alternatives, there are \\emph{non-calorimetric} models that do account for a possible CRE escape \\citep{helou_93a,niklas_97a,lacki_10a,dumas_11a,tabatabaei_13a}. The models of \\citet{helou_93a} and \\citet{lacki_10a} were proposed to explain the RC--far-IR (FIR) correlation, allowing for a possible escape of dust heating FUV photons to compensate the escape of CREs. These models are hence not strictly applicable to the RC--SFR relation, because, as we will see, we are using SF tracers that do account for the escape of FUV photons. Hence, we do not consider them any further. The model by \\citet{niklas_97a} and their spatially resolved derivatives \\citep{dumas_11a,tabatabaei_13a}, however, although again posed to explain the RC--FIR correlation, do assume a direct proportionality between the FIR luminosity and the SFR. Hence, they are essentially models for the RC--SFR relation and thus relevant for our study.  We will use Condon's relation as the basis of our study and investigate whether it holds for our sample for integrated SFRs, as well as for SFR surface densities ($\\Sigma_{\\rm SFR}$). Two emission processes contribute to the RC emission: massive stars contribute to both the thermal and the non-thermal emission. The thermal emission originates from free--free emission of electrons in the \\ion{H}{2} region surrounding these stars. The H$\\alpha$ emission can be used as a proxy for the thermal emission; the radio spectral index of the thermal emission is flat with $\\alpha=-0.1$ ($S_\\nu\\propto \\nu^\\alpha$).  The non-thermal emission is generated by young CREs accelerated in supernova shockwaves. They emit synchrotron emission in the magnetized interstellar medium with a non-thermal spectral index of typically $\\alpha_{\\rm nt}= -0.83$ \\citep{niklas_97b}. \\citet{tabatabaei_07c,tabatabaei_07b} separated the thermal and non-thermal RC emission in M33 based on a novel method to correct Balmer H$\\alpha$ emission for internal extinction by dust. They found that the non-thermal emission is distributed more smoothly than the thermal emission, as expected from cosmic-ray transport \\citep{tabatabaei_07a}.  Young supernova remnants are visible as strong non-thermal sources but contribute only about $10\\%$ to the instantaneous non-thermal RC emission at wavelengths around $\\lambda$20~cm \\citep{condon_92a,lisenfeld_00a}.  The main contribution to the RC emission comes from diffuse emission in the interstellar medium. Cosmic rays diffuse along magnetic field lines or are advected in galactic outflows. The electrons lose their energy and remain visible typically for as long as $10^8~\\rm yr$ after the SF event that generated them, depending on the magnetic field strength (for a strong field, this time can be shorter). The non-thermal RC distribution can therefore be understood as a smeared-out version of the SFR distribution \\citep{bicay_89a,murphy_06a,murphy_06b,murphy_08a}. As we can see, the thermal RC and non-thermal RC behave in different ways, motivating a separation of the RC emission in order to study them separately. Even though the non-thermal RC dominates at wavelengths around $\\lambda$20~cm, the thermal RC can become more important for local measurements (1~kpc scale). In the following, RC refers to the \\emph{total} RC, i.e., the sum of the thermal and non-thermal RC, except when we specifically refer to one of the components.  Studies of the RC--SFR relation started out based on integrated properties of the galaxies under study. Often, no allowance could be made for contaminating background sources within the fairly large beams of single-dish telescopes. Likewise, single-dish radio flux densities can be contaminated by the presence of an active galactic nucleus (AGN. \\citet{condon_02a} tried to separate AGNs from their sample by taking out the outliers with too high a radio luminosity compared to the SFR observed, but that selection is obviously not able to detect some of the weaker AGNs. Spatially resolved studies circumvent this problem as an AGN can be readily identified as an excess of radio brightness in the center of a galaxy. Similarly, background radio sources can be identified and their contribution removed.  Recent observational efforts have concentrated on studies of the RC--FIR or RC--mid-infrared (RC--MIR) correlation rather than the RC--SFR relation at ever higher spatial resolutions. Using wavelet correlation spectra, \\citet{dumas_11a} show that the correlation between $\\lambda 20~\\rm cm$ RC and $24~\\rm\\mu m$ MIR emission in M51 (NGC~5194) holds well down to a scale of $1.8~\\rm kpc$ (the typical width of a spiral arm) and to a less extent to a scale of 500~pc (the CRE diffusion length), but breaks down at smaller scales. A study of the spatially resolved RC--FIR correlation in the Small Magellanic Cloud by \\citet{leverenz_13a} found a good correlation down to a scale of 15~pc; however, this galaxy has a high thermal RC fraction, meaning that cosmic-ray diffusion is of less importance.  Other papers dealing with spatially resolved studies of the RC--FIR correlation are those by \\citet{tabatabaei_07a} on M33 and \\citet{tabatabaei_13a} on NGC~6946, among others.  \\begin{deluxetable*}{lrrrrrrrrrrrr} \\tabletypesize{\\scriptsize} \\tablecaption{Basic Data of the Galaxies in Our Sample.\\label{tab:data_sample}} \\tablewidth{0pt} \\tablehead{ \\colhead{Galaxy\\B} & \\colhead{FWHM} & \\colhead{$i$} & \\colhead{$D$} & \\colhead{$D_{25}$} & \\colhead{$M_{\\rm B}$} & \\colhead{$M_{\\rm HI}$}  & \\colhead{$M_{\\rm H2}$} & \\colhead{$\\log_{10}(M_{\\rm s})$} & \\colhead{$PI$} & \\colhead{Type} & \\colhead{$z$} & \\colhead{$Z$} \\\\ & & & \\colhead{(Mpc)} & & & \\colhead{($10^8M_\\odot$)} & \\colhead{($10^8M_\\odot$)} & \\colhead{($M_\\odot$)} & \\colhead{(\\%)} & & $(10^{-2})$ &  \\\\ \\colhead{(1\\T)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} & \\colhead{(11)} & \\colhead{(12)} & \\colhead{(13)}} \\startdata Ho~II     & $13\\farcs 5$ & $41\\degr$ & $3.4$ & $6\\farcm 61$  & $-16.87$ & $6.0$   & $<0.4$ & $8.3$  & $<1.8$ & 10 & $0.05$ & $7.68$ \\\\ IC~2574   & $13\\farcs 5$ & $53\\degr$ & $4.0$ & $12\\farcm 88$ & $-18.11$ & $14.8$  & $<0.8$ & $8.7$  & $<9  $ & 9  & $0.02$ & $7.94$ \\\\ NGC~628   & $46\\farcs 0$ &  $7\\degr$ & $7.3$ & $9\\farcm 77$  & $-19.97$ & $38.0$  & $10.0$ & $10.1$ & $12  $ & 5  & $0.22$ & $8.33$ \\\\ NGC~925   & $23\\farcs 0$ & $66\\degr$ & $9.2$ & $10\\farcm 72$ & $-20.04$ & $45.8$  & $2.5 $ & $9.9$  & $<0.8$ & 7  & $0.18$ & $8.24$ \\\\ NGC~2403  & $13\\farcs 6$ & $63\\degr$ & $3.2$ & $15\\farcm 85$ & $-19.43$ & $25.8$  & $0.2 $ & $9.7$  & $7.8 $ & 6  & $0.04$ & $8.31$ \\\\ NGC~2841  & $16\\farcs 1$ & $73\\degr$ & $14.1$& $6\\farcm 92$  & $-21.21$ & $85.8$  & $3.2 $ & $10.9$ & $6.2 $ & 3  & $0.21$ & $8.52$ \\\\ NGC~2903  & $34\\farcs 1$ & $65\\degr$ & $8.9$ & $11\\farcm 75$ & $-20.93$ & $43.5$  & $21.8$ & $10.0$ & $4.0 $ & 4  & $0.19$ & $9.12$ \\\\ NGC~2976  & $13\\farcs 3$ & $65\\degr$ & $3.6$ & $7\\farcm 24$  & $-17.78$ & $1.4$   & $0.6 $ & $9.1$  & $4.4 $ & 5  & $0.00$ & $8.30$ \\\\ NGC~3184  & $18\\farcs 6$ & $16\\degr$ & $11.1$& $7\\farcm 41$  & $-19.92$ & $30.7$  & $15.9$ & $10.3$ & $6.4 $ & 6  & $0.20$ & $8.48$ \\\\ NGC~3198  & $17\\farcs 6$ & $72\\degr$ & $13.8$& $6\\farcm 46$  & $-20.75$ & $101.7$ & $6.3 $ & $10.1$ & $<1.4$ & 5  & $0.22$ & $8.32$ \\\\ NGC~3627  & $55\\farcs 6$ & $62\\degr$ & $9.3$ & $10\\farcm 23$ & $-20.74$ & $8.2$   & $12.6$ & $10.6$ & $2.5 $ & 3  & $0.24$ & $8.43$ \\\\ NGC~4736  & $19\\farcs 1$ & $41\\degr$ & $4.7$ & $7\\farcm 76$  & $-19.80 $& $4.0$   & $4.0 $ & $10.3$ & $4.3 $ & 2  & $0.10$ & $8.31$ \\\\ NGC~4826  & $33\\farcs 6$ & $65\\degr$ & $7.5$ & $10\\farcm 47$ & $-20.63$ & $5.5$   & $18.1$ & $10.4$ & $1.0 $ & 2  & $0.14$ & $8.59$ \\\\ NGC~5055  & $18\\farcs 6$ & $59\\degr$ & $10.1$& $11\\farcm 75$ & $-21.12$ & $91.0$  & $50.1$ & $10.8$ & $4.7 $ & 4  & $0.17$ & $8.42$ \\\\ NGC~5194  & $17\\farcs 1$ & $42\\degr$ & $8.0$ & $7\\farcm 76$  & $-21.04$ & $25.4$  & $25.1$ & $10.6$ & $6.9 $ & 4  & $0.15$ & $8.54$ \\\\ NGC~6946  & $14\\farcs 1$ & $33\\degr$ & $5.9$ & $11\\farcm 48$ & $-20.61$ & $41.5$  & $39.8$ & $10.5$ & $10  $ & 6  & $0.02$ & $8.40$ \\\\ NGC~7331  & $22\\farcs 1$ & $76\\degr$ & $14.7$& $9\\farcm 12$  & $-21.67$ & $91.3$  & $50.1$ & $10.9$ & $3.8 $ & 3  & $0.27$ & $8.36$ \\enddata \\tablecomments{Column 1: galaxy name; Column 2: resolution of the radio map as full-width-half-mean in arcsec; Column 3: inclination angle in degree; Column 4: distances in Mpc; Column 5: optical size in arcmin; Column 6: $B$-band absolute magnitude from \\citet{walter_08a}; Column 7: hydrogen mass from \\citet{walter_08a}; Column 8: molecular hydrogen mass from \\citet{leroy_08a}, NGC~2903 is from \\citet{leroy_09a}, NGC~4826 from \\citet{wilson_12a}; Column 9: decadal logarithm of stellar mass from \\citet{leroy_08a}, NGC~2841 and NGC~2903 are from \\citet{deblok_08a}, NGC~4826 was determined from \\emph{IRAS} data as in \\citet{leroy_08a}; Column 10: fractional polarization at $\\lambda$22~cm from \\citet{heald_09a} with respect to our integrated flux density; Column 11: galaxy type from \\citet{walter_08a}; Column 12: redshift; Column 13: metallicity defined by $\\log_{10}({\\rm O/H})+12$ from \\citet{moustakas_06a}.} \\end{deluxetable*}  But until recently, no studies of the RC--SFR have been forthcoming. Obviously, such studies should have sufficient angular resolution in order to be able to isolate any AGN activity or unrelated background sources from the area under study. Moreover, spatially resolved studies allow one to probe the RC--SFR relation as a function of location within a galaxy, such as arm versus inter-arm. Thanks to the THINGS \\citep[The \\ion{H}{1} Nearby Galaxy Survey,][]{walter_08a} collaboration---making use of SINGS \\citep[SIRTF Nearby Galaxy Survey;][]{kennicutt_03a} and the ``GALEX Ultraviolet Atlas of Nearby Galaxies'' \\citep{gildepaz_07a}---spatially resolved SFR maps have become available of a significant sample of nearby galaxies. In this paper, we present a sample of 17 galaxies where we compare the RC brightness and the SFR surface density ($\\Sigma_{\\rm SFR}$). We use $\\Sigma_{\\rm SFR}$ maps based on a combination of \\emph{Spitzer} MIR and \\emph{GALEX} FUV maps \\citep[see][for details]{leroy_08a}. The RC data are from the Westerbork Synthesis Radio Telescope (WSRT) SINGS sample \\citep{braun_07a}, taken at $\\lambda\\lambda$ 22 and 18~cm.  The paper is organized as follows: Section~\\ref{sec:observations} explains the selection of the sample and briefly summarizes the data processing applied to the maps that we used. Section~\\ref{sec:methodology} deals among other things with the azimuthally (Section~\\ref{subsec:radial}) averaged plots where we compare the RC brightness with the $\\Sigma_{\\rm SFR}$ as a function of galactocentric radius. In Section~\\ref{subsec:pixel} we produce pixel-by-pixel plots of our sample at a fixed resolution. The results are analyzed in Section~\\ref{sec:results} and discussed in Section~\\ref{sec:discussion} where we explore the dependence of the RC--SFR relation on various galaxy parameters. We present our conclusions in Section~\\ref{sec:conclusions}. ",
        "conclusions": ""
    },
    "1402/1402.6900_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:introduction} \\setcounter{footnote}{0} \\setcounter{equation}{0}  Our universe is well described by the so-called spatially-flat $\\Lambda$CDM model. Energy density of our universe is close to the critical density. According to the PLANCK 2013 results~\\cite{PlanckCosm}, only 5.1\\% of the energy density is attributed to a known form of baryonic matter, while 26.8\\% is attributed to cold dark matter (weakly interacting non-relativistic matter), and 68.3\\% to dark energy (or the cosmological constant). Explaining the origin of these unknown ingredients, dark matter and dark energy, is one of the biggest challenges in modern physics.  There are various proposals for dark matter, such as supersymmetric particles, axions, and so on. Dark matter may well be one of these. However, the origin of dark energy is totally unclear. Although its equation of state, $w(=p/\\rho)=-1$, seems like that of vacuum energy of quantum fields, there is no reasonable explanation for its magnitude, $\\rho_{\\rm{DE}}= 3 (M_P H_0)^{2} \\Omega_\\Lambda \\sim (2.2{\\rm meV})^4$ with meV=$10^{-3}$eV. Here, $M_P = (8 \\pi G_N)^{-1/2} \\sim 2.4\\times 10^{30} {\\rm meV}$ is the (reduced) Planck scale and $H_0 \\sim 1.4\\times 10^{-30} {\\rm meV}$ is the current Hubble parameter. $\\rho_{\\rm{DE}}$ is far smaller than the expected magnitude of the vacuum energy $m^4$ in a theory with an ultra-violet (UV) cutoff, $m$, with any reasonable choice for $m$. If we take $m$ to be $M_P$, $\\rho_{\\rm{DE}}$ is smaller than $m^4$ by more than 120 orders of magnitude. Even if we take $m$ to be the supersymmetry breaking scale, electroweak scale, or any other natural scale in high-energy physics, $\\rho_{\\rm{DE}}$ is still much smaller than $m^4$. This is the cosmological constant problem~\\cite{WeinbergCC}. It may turn out that the solution to this problem is  given by the anthropic principle~\\cite{WeinbergAP}, but attempts at a dynamical explanation of dark energy are undoubtedly important.  In this paper, we study some aspects of vacuum energy of quantum fields in cosmological spacetimes, with an application to the cosmological constant problem in mind. Let us first note that the order of magnitude of the energy density of quantum fields is not necessarily given by the value of the UV cutoff. Expectation values of the energy-momentum tensor should be renormalized by subtracting the cutoff-dependent (divergent) term. There is a well-defined method for regularization and renormalization, which yields finite expectation values for energy-momentum tensors, which are covariant and conserved~\\cite{Birrell:1982ix}. The terms to be subtracted are a combination of spacetime curvature tensors. Typical terms in the renormalized expectation value (such as the terms responsible for the Weyl anomaly) are of the order of the background curvature.  In particular, vacuum energy for fluctuations in the de Sitter background with the Hubble parameter $H$ is of the order $H^4$ (for massless fields), and has $w=-1$~\\cite{BunchDavies, Dowker}. Since our present universe is close to de Sitter space, one may wonder if dark energy can be explained as vacuum energy in de Sitter with the current Hubble parameter $H_{0}$, but this does not seem to be plausible. Dark energy that we observe, $M_P^2H_{0}^{2}$, is much larger than the expected contribution from a single field, $H_{0}^4$.  However, local curvature is not the only dimensionful quantity that affects the renormalized energy-momentum tensor. To compute expectation values of fluctuations, we need to specify the vacuum state. This may depend on global properties of the geometry and the whole history of the universe; thus, different scales might be introduced in the problem.  There is by now strong evidence~\\cite{PlanckInf} that there has been a period of inflation in our past. The spacetime during inflation is nearly de Sitter space with the Hubble parameter $H_I$ being much larger than $H_0$. Inflation should have lasted long enough (e-foldings $N_{e} \\gtrsim  60$) to solve the flatness and horizon problems. It would be reasonable to take the vacuum to be the Bunch-Davies vacuum~\\cite{BunchDavies} for de Sitter space with $H_I$. The Bunch-Davies vacuum is the one obtained by the Euclidean prescription, and reduces to the Minkowski vacuum in the short wavelength limit. Even if different initial states are taken, correlation functions will be attracted to those taken with respect to the Bunch-Davies vacuum, as shown quite generally \\cite{HabibMottola}.  Fluctuations of the massless scalar in the de Sitter background (in four spacetime dimensions) is of the order $H_I$, as is clear from dimensional analysis. Massless fluctuations are frozen (remain constant) outside the Hubble radius (see, {\\it e.g.}, \\cite{Mukhanov:2005sc}); thus, infra-red (IR) modes could have a large value in the universe after inflation. In fact, these fluctuations are considered to be the origin of the fluctuations of the cosmic microwave background (CMB) that is observed today~\\cite{PlanckCosm, PlanckInf}.  The purpose of this paper is to understand how the IR modes of quantum fields could affect the energy density in the present universe. We will mostly consider a massless minimally-coupled scalar field, and study the time evolution of the energy-momentum tensor in detail. Our result will shed light on the effect of almost massless and non-interacting fields, such as axions. Our work is related to the studies of the fluctuations of the graviton or the inflaton~\\cite{Abramo1, Abramo2, Abramo3}, but further modifications are necessary since energy-momentum tensors for gravitons have different tensor structures from those for scalars and those for the inflaton have contributions from the classical value of the field. We believe our work serves as a starting point for the study of the time evolution of those fields.\\footnote{ There is a paper \\cite{Brown:2009cy}, in which the authors claim that back reactions of gravitational waves produced during inflation can serve as a source for the acceleration of the present universe, by taking an approach somewhat different from ours. }  We will find that the magnitude of the present energy-momentum tensor of a massless field is of the order $H_I^{2} H_0^{2}$, with a possible factor logarithmic in the scale factor. The contribution from each momentum mode $k$ can be renormalized separately before integration over $k$ is performed. The equation of state for the (un-renormalized) contribution from each $k$ has $w>-1/3$: The IR mode $k\\to 0$ has $w\\to -1/3$; for larger $k$, $w$ is larger, and the UV limit has $w\\to 1/3$, which is the same as radiation. The equation of state may become $w<-1/3$ due to renormalization, but the terms arising from renormalization will be of the order of the background curvature. In the present universe, these will be of the order $H_0^4$, and are negligible.\\footnote{ After submitting our paper, we became aware of a recent paper~\\cite{Glavan:2013mra}, which has substantial overlap with the first part of our paper. In \\cite{Glavan:2013mra}, the authors study vacuum fluctuations of a massless minimally coupled test field in a background that evolves from inflation to radiation domination to matter domination. They use the approximation of sharp transition between these eras, just like we do. The results in the first part of our paper (which is summarized in this paragraph) is completely consistent with the ones in \\cite{Glavan:2013mra}. }  The value $H_I^{2} H_0^{2}$ is smaller than dark energy (or total energy density) in our universe, since the scale of inflation is smaller than the Planck scale, $H_{I}/M_P<3.6 \\times 10^{-5}$, as suggested by the observations of the CMB~\\cite{PlanckInf}. But, in this paper, we point out a possibility that the energy density of a quantum field takes a larger value.\\footnote{This is different from proposals based on secular growth of fluctuations in inflation~\\cite{Kolb, Barausse}. We are not assuming the expectation value of energy density in pure de Sitter grows with time.} Even though the e-folding of inflation has to be large enough to make the region inside our horizon smooth, it does not have to be infinite. Finiteness of the e-foldings means the long wavelength modes have not been in causal contact during the inflationary era. There is no reason to expect that those modes are in the Bunch-Davies vacuum. It would not be too surprising if there are large fluctuations in the far IR. After all, the point of inflation is to push away inhomogeneities beyond our horizon. In particular, in the context of eternal inflation~\\cite{CDL, HM, Linde}, our universe is generally surrounded by the region with a larger Hubble parameter where there are large fluctuations.  It is difficult to know what happened before inflation. In this paper, as an explicit example, we consider a double inflation model. We assume there was an inflation with the Hubble parameter $H_P$ of the order of the Planck scale $M_P$ (such as Starobinsky's inflation~\\cite{Starobinsky}, for example), followed by either a radiation-dominated or curvature-dominated transition period, before the usual inflation with $H_I$ starts. We fix the initial condition of the fields in the Planckian inflation period by taking the Bunch-Davies vacuum for de Sitter with the Hubble $H_P$, and study the time evolution afterwards. In this case, IR mode is enhanced to $H_P$.  By computing the evolution of the energy-momentum tensor for the double inflation model, we find the present value of vacuum energy to be of the order of $H_{P}^{2} H_{0}^{2}$. One may worry that the fluctuations become so large that they contradict the observed value of CMB fluctuations. This will depend on what field we are considering, and needs careful study. In this paper, we argue that it is possible to enhance the modes that have a longer wavelength than the scales that are observed in CMB, leaving shorter wavelength modes essentially unaffected.  The message of this paper is that vacuum energy of the order of the energy density of our present universe may arise, due to the enhancement of the IR fluctuations generated in the very early universe. In our analysis of free fields, we were able to obtain only $w>-1/3$, which cannot drive acceleration. However, we should mention that the free field approximation is not likely to be valid in the far IR where large fluctuations generated before the ordinary inflation exist. At the end of the paper, we will mention possible directions for future study to take interactions into account.  The paper is organized as follows. We review the basics of quantization of scalar fields in curved spacetimes in Section~\\ref{sec:review}. After specifying the cosmic history of background  geometry in Section~\\ref{sec:history}, we obtain the wave function of a massless minimally-coupled scalar field, with the initial condition fixed in the inflationary era in Section~\\ref{sec:wavefunc}. Then we calculate the energy-momentum tensor, paying special attention to the contributions from the IR modes in Section~\\ref{sec:emt}. We explain the prescription for treating the UV divergence, and mention subtleties associated with a  physical interpretation of cutoff depenent terms, and present the renormalized energy-momentum tensor in Section~\\ref{sec:renormalization}. We consider time evolution of the energy-momentum tensor from the inflationary era until the present in Section~\\ref{sec:timeevo}. In Section~\\ref{sec:double}, we consider a double inflation model, and discuss the effects of a period that may have preceded the ordinary inflation. In section \\ref{sec8.2}, we summarize  time evolution of energy densities generated in the inflation period and in the pre-inflation period. They are shown in Figure~\\ref{fig:timeevorhos}. Section~\\ref{sec:discussion} is devoted to conclusions and discussion. In Appendix~\\ref{sec:chasmmodel}, we consider a double inflation model with a different intermediate stage. In Appendix~\\ref{sec:IRbhvwf}, we investigate IR behaviors of the wave functions. ",
        "conclusions": "\\label{sec:discussion} \\setcounter{equation}{0}  In this paper, we calculated the time evolution of the energy-momentum tensor of a minimally-coupled massless scalar field throughout the history of the universe. We considered two types of cosmic histories. The first one is the standard cosmology model, starting from the inflation and followed by the RD and MD periods. To perform quantization, the Bunch-Davies initial condition is imposed on the field in the inflation period. Due to the fact that inflation produces fluctuations of the order $H_I$, the energy density of a quantum field becomes of the order $\\rho_{\\rm inf}\\sim (H_I H)^2$ where $H$ and $H_I$ are the Hubble parameters at each moment and in the inflation period. The evolution of  $\\rho_{\\rm inf}$ is given by eq.~(\\ref{rhoInfevo}) and its ratio to the critical density is depicted as the dashed line in Figure~\\ref{fig:timeevorhos}. The ratio is almost constant but the magnitude is much smaller than the critical density. The equation of state evolves from a negative value to $w=1/3$ in the RD period, and from $1/3$ to $w=0$ in the MD period.  The second type of model we considered is a double inflation model, in which we assume another inflation, which we call a pre-inflation, exists before the ordinary inflation starts. The present energy density of the fluctuations generated in the pre-inflation is enhanced to an order of $\\rho_{\\rm pre-inf} \\sim (H_P H_0)^2$ and is comparable to the critical density if $H_P \\sim M_P$. Here, $H_P$ is the Hubble parameter of the pre-inflation and $M_P$ is the Planck scale. The evolution of $\\rho_{\\rm pre-inf}$, given by eq.~(\\ref{rhoPreInfevo}), is  different from $\\rho_{\\rm inf}$ because the amplified wave functions in the pre-inflation have larger wavelengths and have not entered the horizon yet. The ratio to the critical density is depicted by the solid line in Figure~\\ref{fig:timeevorhos}. The equation of state is given by $w=p/\\rho = -1/3$. It is negative, but cannot drive the acceleration of our universe within the free field approximation.  Our result suggests that fluctuations generated before the ordinary inflation could have an important effect on the present universe. We have considered the double inflation models as a simple explicit example in this paper. There are theoretically well-motivated models in the context of eternal inflation. In a theory with a metastable vacuum, universes are created by bubble nucleation~\\cite{CDL}; such universes are surrounded by de Sitter space (the ancestor vacuum) with the Hubble parameter being supposedly much higher than the one for the ordinary inflation $H_I$. The vacuum state in such a universe can be defined by the Euclidean prescription, and is different from the Bunch-Davies vacuum for ordinary inflation~\\cite{FSSY}. It is an important problem to find the renormalized energy-momentum tensors in such a universe, extending the analysis of this paper.  In this work, we have not considered back reactions from the quantum fluctuations to the geometry. In recent papers~\\cite{Habara1, Habara2}, in which the effect of a large number of fields (Kaluza-Klein and string states) on the CMB fluctuations was studied, the authors made an interesting proposal that the vacuum energy from these fluctuations itself drives the acceleration in inflation. Back reactions to the geometry in this context are being studied by these authors. In our work, the energy density of the fluctuations is much smaller than the critical value during most of the periods in the cosmic history. However, it becomes relevant at the very early universe when the Hubble parameter is close to the Planck scale, and also at times later than the present where the vacuum fluctuation found in this paper starts to be dominant. It would be a very interesting theoretical problem to study the late time behavior to understand the fate of the universe. In this region, the interplay among the scale factor, the IR behavior of the  wave function, and the Bunch-Davies initial condition should determine the dynamics of the universe self-consistently.  We solved the evolution equation in the free field approximation. The system is very simple, yet the evolution of the energy density of a quantum field summarized in Figure~\\ref{fig:timeevorhos} is  nontrivial. Then what will happen if we include  interactions among them? Although there have been debates on the physical effect of quantum loops in the de Sitter background (see, {\\it e.g.}, \\cite{Tanaka} and references therein), it would be reasonable to assume that massive (including light but not exactly massless) fields will reach the interacting Hartle-Hawking vacuum~\\cite{Marolf}. If this is the case, the effect of loops modifies the numerical coefficients in front of $\\rho$, but not their qualitative behaviors.  As a final remark, we mention a possible scenario of how interactions change the equation of state. As is well known in the Bogoliubov theory of superfluidity, quantum Bose gases fall into a macroscopic state in the presence of interactions. If the interaction is attractive, the system is unstable with negative pressure. Let us assume here that condensation occurs in our system. Indeed, the scalar field has gravitational interaction. Since the gravitational interaction is very weak and the energy density is also very low, it will take a long time for the condensation to occur. Once the condensation occurs, the interaction energy dominates the kinetic energy and the equation of state will be changed from the free case considered in the present paper to the  interaction dominated form. Since the gravitational interaction is attractive, the negative pressure is expected to arise. It is interesting to understand  the time scale of the condensation dynamics and evolutions of the energy-momentum tensor modified by gravitational interactions. We will investigate these issues in future publications."
    },
    "1402/1402.4134_arXiv.txt": {
        "abstract": "Massive stars that end their lives with helium cores in the range of 35 to 65 \\Ms\\ are known to produce repeated thermonuclear outbursts due to a recurring pair-instability.  In some of these events, solar masses of material are ejected in repeated outbursts of several times 10$^{50}$ erg each. Collisions between these shells can sometimes produce very luminous transients that are visible from the edge of the observable universe.  Previous 1D studies of these events produce thin, high-density shells as one ejection plows into another.  Here, in the first multi-dimensional simulations of these collisions, we show that the development of a Rayleigh-Taylor instability truncates the growth of the high density spike and drives mixing between the shells.  The progenitor is a 110 \\Ms\\ solar-metallicity star that was shown in earlier work to produce a superluminous supernova.  The light curve of this more realistic model has a peak luminosity and duration that are similar to those of 1D models but a structure that is smoother. ",
        "introduction": "The idea of a ``pulsational pair-instability supernova'' (PPISN) was introduced by \\citet{barkat1967} and explored in some detail by \\citet{woosley2007}. For a range of helium core masses above approximately 35 \\Ms, the production of electron-positron pairs {\\color{black} occurs during central neon burning} and triggers an instability that leads to rapid contraction of the core and explosive nuclear burning.  If the helium core mass is above about 65 \\Ms, {\\color{black} the pair production instability occurs after carbon ignition} and the energy released can completely unbind the star in a single explosive event known as a pair-instability supernova \\citep[PISN;][]{heger2002,scannapieco2005,chat2012b, chat2013, chen2014b}. If the core mass is above 133 \\Ms, for reasonable nuclear reaction rates, convection theory, and no rotation, nuclear burning is unable to reverse the collapse and a black hole forms \\citep{heger2002}.  For helium core masses from 35 - 65 \\Ms,  Although explosive burning is violent and energetic, it cannot unbind the entire star, and cycles of instability and mass ejection can instead occur. The energy, duration, and mass ejected by these nuclear-powered pulses increases as the helium core mass rises, and by about 45 \\Ms\\ they are sufficient to produce SN-like displays {\\color{black} \\citep{woosley2007,vink2014}}. If the helium core is capped by a substantial hydrogen envelope, the first strong pulse ejects it.  Depending on the radius and mass of the envelope, this initial ejection may give rise to either a faint or rather typical Type IIp SN. Subsequent pulses later overtake and collide with the first and produce a much brighter Type IIn SN.  If there is no hydrogen envelope, collisions between helium shells can produce a bright Type I SN.  The helium core mass range from 45 to 55 \\Ms\\ is particularly interesting because the characteristic time scale between pulsations is years and the collisions between shells ejected with speeds $\\sim$ 1000 km s$^{-1}$ occur at 10$^{15}$ - 10$^{16}$ cm, where the collision energy is mostly dissipated by optical emission.  Since the collision energy can approach 10$^{51}$ erg, a superluminous event can result.  The discovery of both PI and PPI SN candidates in the local universe, and the realization that such events might be visible at high redshifts, have excited interest in this exotic explosion mechanism.  SN 2007bi at $z =$ 0.127 \\citep{2007bi} and SN 2213-1745 at $z =$ 2.06 \\citep{cooke12} are PISN candidates, and SN 1000$+$0216 at $z =$ 3.90 \\citep{ cooke12} and perhaps SN 2006oz at $z =$ 0.376 \\citep{lel12} may be PPISNe.  {\\color{black} \\citet{dessart2013} suggest that SN 2007bi may be a magnetar spin-down powered event \\citep{kasen2010, woosley2010, dessart2012} or it could be due to an interaction between H-poor SN ejecta and a circumstellar medium \\citep{chat2012a, moriya2013}. The nature of these transients is still under debate.} \\citet{wet13d} have now shown that PPISNe, like PISNe and Type IIne, may be visible to the {\\it James Webb Space Telescope} ({\\it JWST}) and the Wide-Field Infrared Survey Telescope (WFIRST) at $z \\sim$ 20 \\citep[see also][]{ kasen2011,pan12a,wet12b,wet12a,wet12e,hum12}. These explosions, together with other Population III (Pop III) SNe \\citep{candace2010,wet12c,wet12d,jet13a}, could probe the properties of the first stars and galaxies \\citep{fsg09,glov12,dw12,fg11}, early cosmological reionization and chemical enrichment \\citep{whalen2004,mbh03,ss07,bsmith09,ritt12,ss13}, and the origins of supermassive black holes \\citep{vol12}.  Understanding the observational signatures of PPISNe is key to properly identifying them as more of them are discovered by the new SN factories such as the Palomar Transient Factory \\citep[PTF;][]{ptf1}, the Panoramic Survey Telescope and Rapid Response System \\citep[Pan-STARRS;][]{panstarrs} and the Large Synoptic Survey Telescope \\citep[LSST;][]{ lsst}.  High-$z$ SN surveys by {\\it JWST} and WFIRST may harvest even greater numbers of PPISNe if, as many suspect, the Pop III initial mass function (IMF) is top-heavy.  \\citet{ woosley2007} modeled a PPISN in one dimension (1D) with the \\KEPLER{} code \\citep{kepler, heger2001} and its light curve with the \\STELLA{} code \\citep{stella}.  In these simulations, {\\color{black} a large density spike formed during the collision between the second two pulses and the first.}  Rapid variations in the density of the spike due to a radiative instability \\citep{ chev82,imam84} imposed large fluctuations in the bolometric luminosity. \\citet{wet13d} noted similar features in the near infrared (NIR) light curves of the same PPISN modeled with the Los Alamos \\RAGE{} code. Such spikes would probably not appear in multidimensional flows because hydrodynamic instabilities would likely erase the sharp interface between the pulses. It is not clear how these processes would change the luminosity of the collision.  Radiation hydrodynamical simulations in at least two dimensions (2D) with adaptive mesh refinement (AMR) are clearly required to resolve the thin radiating region between the shells during the collision and properly model the light curves of PPISNe.  But such simulations push the envelope of even state-of-the-art numerical codes.  As a first step to this goal, we have performed the first 2D simulations of a PPISN with the \\CASTRO{} code with hydrodynamics but not radiation transport.  In Section 2 we describe our PPISN model and how it is evolved in \\CASTRO.  In Section 3 we examine the collision of the shells in 2D in detail, and in Section 4 we show how the evolution of the collision in 2D might change the PPISN light curves of previous calculations.  We conclude in Section 5. ",
        "conclusions": "We have performed the first 2D hydrodynamical simulations of PPISNe with the new \\CASTRO{} code.  We find that mild RT instabilities develop first during the collision of P3 with P2 and then become stronger as P2 and P3 interact with P1.  Fallback from one pulsation onto another does not seed strong instabilities but does mildly perturb the flow prior to collisions.  The appearance of dynamical instabilities and mixing has three consequences for the observational signatures of PPISNe that are not captured in 1D models.  First, the radiative instabilities that cause light curves to fluctuate in 1D models probably do not happen in actual explosions.  Mixing between shells dampens the formation of the density spikes in 1D models that give rise to these features.  Second, the formation of dense clumps by RT fingers could trap photons and alter the luminosity of these events.  We neglect radiative cooling by metals and dust in our models so these clumps may catastrophically cool and become much denser (and more opaque) in actual collisions.  This in turn could amplify the RT instabilities. Finally, since mixing is strongest in the region of the flow from which most of the luminosity of the collision originates, it will alter the order in which specific absorption and emission lines appear in the spectrum over time.  In particular, our model predicts enhanced spectral lines from \\Cx{} and \\Ox{} at earlier times in the second collision.  The appearance of such lines may prove to be powerful diagnostics of mixing in PPISNe.  Since our approximate bolometric luminosities are similar to those calculated with full radiation transport in earlier work, future multidimensional calculations will probably not change the recent result that these explosions will be visible at $z =$ 10 - 20 to {\\it JWST} and WFIRST.  But they do demonstrate that multidimensional radiation transport will be required to model how radiation is emitted by and escapes from the complex structures formed by dynamical instabilities.  This will be key to obtaining more realistic light curves and spectra for PPISNe.  In future papers, we will better survey PPISN explosions by examining more progenitor models at low metallicities with 2D radiation hydrodynamical calculations \\citep{zhang2013}.  PPISNe will soon open new windows on massive star formation in both the primordial and the local universe."
    },
    "1402/1402.1182_arXiv.txt": {
        "abstract": "Mixing length theory is the predominant treatment of convection in stellar models today.  Usually described by a single free parameter, $\\alpha$, the common practice is to calibrate it using the properties of the Sun, and apply it to all other stellar models as well.  Asteroseismic data from \\textit{Kepler} and \\textit{CoRoT} provide precise properties of other stars which can be used to determine $\\alpha$ as well, and a recent study of stars in the \\textit{Kepler} field of view found $\\alpha$ to vary with metallicity.  \\rev{Interpreting $\\alpha$ obtained from calibrated stellar models}, however, is complicated by the fact that the value for $\\alpha$ depends on the surface boundary condition of the stellar model, or {\\ttau} relation.  Calibrated models that use typical {\\ttau} relations, which are static and insensitive to chemical composition, do not include the complete effect of metallicity on $\\alpha$.  We use 3D radiation-hydrodynamic simulations to extract metallicity-dependent {\\ttau} relations and use them in calibrated stellar models.  We find the \\rev{previously reported} $\\alpha$-metallicity trend to be robust, and not significantly affected by the surface boundary condition of the stellar models. ",
        "introduction": "\\label{sec:introduction}  The Mixing Length Theory \\citep[MLT;][]{1958ZA.....46..108B} remains one of the most popular treatments for stellar convection, and typically describes convective eddy sizes with a single parameter, $\\alpha$, which is arbitrarily set by the modeler.  Generally, $\\alpha$ is kept fixed at a value that is determined by calibrating \\rev{a solar model} to the properties of the Sun, however, this  has no \\textit{a priori} justification.  With space-based photometric data from \\textit{CoRoT} \\citep{corot} and \\textit{Kepler} \\citep{kepler}, it is now possible to precisely constrain the properties of other stars, \\rev{which allows $\\alpha$ to be determined in a similar manner.  Determining} $\\alpha$ is possible with tight constraints on the stellar mass and radius, however, it still  depends on a variety of input microphysics and boundary conditions.  It is increasingly apparent that the usual approach of using a constant solar-calibrated value for $\\alpha$ is not appropriate for stars that differ from the Sun, either in composition, mass, or stage of evolution.  As an example, asteroseismic studies of $\\alpha$ Cen \\citep[e.g.][]{dem86,fer95,mig05} require a non-solar value of the mixing length parameter to reproduce the stellar radius.  Stars from \\textit{CoRoT} and \\textit{Kepler} also require non-solar mixing length values to model the oscillation spectra \\citep[e.g.,][]{met10, deh11, mat12}.  More recently, \\citet{2012ApJ...755L..12B} found a systematic metallicity dependence of $\\alpha$ with calibrated stellar models of stars in the Kepler field of view.  It is challenging to interpret the meaning of $\\alpha$ in calibrated stellar models, since it is intertwined with the atmospheric boundary condition and other input microphysics.   The purpose of this study is to test whether the trends reported by \\citeauthor{2012ApJ...755L..12B} are sensitive to the treatment of the surface boundary condition, which is expected to vary with metallicity. \\rev{Note that in the present work, we do not provide a calibration of the mixing length parameter, but rather we investigate how a metallicity-dependent surface boundary condition would affect such a calibration.}  Adjusting the value of $\\alpha$ determines the specific entropy of the convection zone, which in turn sets the stellar radius.  In effect, $\\alpha$ is a free parameter that defines the radius of the stellar model.  The effect of the MLT on the stellar model is most important near the surface, where the specific entropy of the convection zone is set.  This so-called superadiabatic layer (SAL) is several scale heights above and below the stellar photosphere, defined in stellar models as where $T=T_{\\rm eff}$.  Stellar models usually separate the atmosphere (defined as the layers above the {\\teff} surface) from the stellar envelope, and provide an atmospheric structure from a {\\ttau} relation \\rev{which is integrated inward from a very small optical depth}.    This relation, along with the assumption of hydrostatic equilibrium, defines the structure of the outermost layers of the stellar model.  In the \\citeauthor{2012ApJ...755L..12B} study, all of the models were computed with a fixed {\\ttau} relation, which is the usual approach.  However, since the surface boundary condition is determined in part by the nature of the near-surface stellar convection, the relationship between the value of the mixing length parameter and metallicity may be affected by the arbitrary choice of {\\ttau} relation.  In the following sections we describe several tests conducted using stellar models computed with a variety of surface boundary conditions;  these range from the usual prescribed static {\\ttau} relations to using atmospheric stratifications extracted from 3D simulations. ",
        "conclusions": "\\label{sec:trends}   To determine whether the stellar model surface boundary condition significantly affects the metallicity-mixing length trend reported by \\citet{2012ApJ...755L..12B}, we compute several groups of calibrated models similar to those presented in Figure \\ref{fig:calibrated_models} but with different surface boundary conditions.  We make a $0.80 M_\\odot$ cut in the set of models (presented in Fig. \\ref{fig:allm08models}) to show the metallicity-mixing length trend.  The trend is clearly visible, with all low-Z models having smaller mixing length values than those with higher-Z.  Changing the {\\ttau} relation from Eddington to KS introduces a shift in the mixing length value, but leaves the trend almost unchanged.  Introducing simulated {\\ttau} relations shifts the mixing length values as a function of metallicity, and the effect is largest at low-Z.  Models that used simulated surface boundary conditions relations are quite similar to the Eddington models because both {\\ttau} relations are similar at the stellar photosphere.  As mentioned in the introduction, the value of the mixing length parameter is sensitive to the atmospheric boundary condition of the stellar models.  To illustrate this effect we compare the set of Eddington {\\ttau} calibrated models with those that have the KS {\\ttau} relation.  The change in the surface boundary condition results in a shift in the mixing length parameters, shown with dotted lines in Fig. \\ref{fig:allm08models}.     Switching the surface boundary condition between the Eddington and KS {\\ttau} relations, however, does not change the basic behavior as a function of effective temperature or metallicity.   \\begin{figure}[h] \\epsscale{1.15} \\plotone{./m08_all.eps} \\caption{Calibrated stellar models with a stellar mass of 0.80$M_\\odot$ for several {\\ttau} relations.  The metallicity-mixing length trend is clear and independent of the choice of {\\ttau} relation.  \\rev{The effect of introducing a metallicity-dependent {\\ttau} relation is not large enough to change the overall $\\alpha$-metallicity trend, although the relative effect on the mixing length parameter is larger at low-$Z$. }} \\label{fig:allm08models} \\end{figure}   \\citet{2012ApJ...755L..12B} performed a trilinear fit to their data to measure the variation of  the mixing length value as a function of {\\logg}, {\\logteff}, and [Fe/H].  Since our simulated data points share the same {\\logg=4.30}, we perform a similar bilinear fit. \\begin{equation} \\alpha = a + b \\log T_{\\rm eff} + c[Fe/H], \\end{equation} where we define the metallicity as: \\begin{equation} [Fe/H] = \\log(Z/X) - \\log(Z/X)_\\odot. \\end{equation}  Models were weighted for the linear regression such that the models comprising each metallicity group sum to equal values.  The result of the bilinear fit shows a robust correlation between metallicity and the mixing length value, regardless of whether a constant {\\ttau} relation is used, or one that is a function of metallicity and derived from simulations.  The proportionality factor is $c=0.31 \\pm 0.048$ for the complete sample, and $c=0.33 \\pm 0.029$ for the age-restricted subset of models.  In the \\citet{2012ApJ...755L..12B} study, the stellar masses and radii were asteroseismically determined.  In our set of models, it is possible that there is a $\\alpha$-dependence on stellar mass, since our models span a range of stellar masses for a given {\\logg} and {\\logteff}.  In addition to the bilinear fit described above, we also performed a trilinear fit of: \\begin{equation} \\alpha = a + b \\log T_{\\rm eff} + c[Fe/H] + dM/M_\\odot . \\end{equation} The proportionality constant from this fit is $c= 0.74 \\pm 0.019$ and $c= 0.64 \\pm 0.030$ for the complete sample and  age-restricted subset, respectively.  The results of our bilinear and trilinear fitting are robust, and indicate that the relationship between metallicity and the mixing length value, as reported by \\citet{2012ApJ...755L..12B}, are not an artifact of the stellar model surface boundary condition.  This strengthens the understanding that the solar-calibrated value for the mixing length parameter is not suitable for modeling other stars.  Although the metallicity-mixing length trend is significant, it is not possible to directly interpret the mixing length value, as it depends on the details of the formulation of MLT in the stellar model as well as on any input physics that affect entropy, of which the {\\ttau} relation is just one.  For a given MLT formulation and set of input physics, changes in the mixing length parameter are still informative.  We have carried out our tests on stellar models within the mixing length framework because MLT-like treatments of convection are predominant in stellar modeling today.   It is important to note, however, that the MLT formalism and the tunable parameter associated with it is just a proxy for describing convection in stars.  Even if a value for the mixing length can be extracted directly from 3D simulations \\citep[e.g.][]{1999A&A...346..111L, 2011ApJ...731...78T}, MLT cannot correctly reproduce the properties of 3D simulations in the superadiabatic layer near the stellar surface.  Going beyond the surface boundary condition to include other aspects of stellar convection may change the result as well.  This work focuses on the {\\ttau} relation, but it contains only a small part of the information present in 3D simulations.  Including a representation of additional physical processes, such as the turbulent pressure contribution to hydrostatic equilibrium, will potentially yield different results.   It is desirable to ultimately replace MLT-like treatments of convection with a better description of stellar convection (for example, the ongoing efforts of \\citealp{2013ApJ...769....1V}), and such models may behave quite differently than our current MLT models."
    },
    "1402/1402.6594_arXiv.txt": {
        "abstract": "We reconsider the entropic-force model in which both kind of Hubble terms ${\\dot H}$ and $H^{2}$ appear in the effective dark energy (DE) density affecting the evolution of the main cosmological functions, namely the scale factor, deceleration parameter, matter density and growth of linear matter perturbations. However, we find that the entropic-force model is not viable at the background and perturbation levels due to the fact that the entropic formulation does not add a constant term in the Friedmann equations. On the other hand, if on mere phenomenological grounds we replace the ${\\dot H}$ dependence of the effective DE density with a linear term $H$ without including a constant additive term, we find that the transition from deceleration to acceleration becomes possible but the recent structure formation data \\newtext{strongly disfavors} this cosmological scenario. Finally, we briefly compare the entropic-force models with some related DE models (based on dynamical vacuum energy) which overcome these difficulties and are compatible with the present observations. ",
        "introduction": "\\frac{\\ddot{a}}{a}=-\\frac{4\\pi\\,G}{3}\\,(1+3\\wm)\\,\\rmr+\\CHd\\,\\Hd\\,+\\CH\\,H^2 \\end{equation} where $\\wm=\\pmr/\\rmr$ is the equation of state (EoS) for the matter fluid (with $\\wm=0$ and $1/3$ for non-relativistic and relativistic matter, respectively). The second cosmological equation is the generalized Friedmann's equation: \\begin{equation}\\label{GeneralizedFriedmann} H^2\\equiv\\left(\\frac{\\dot{a}}{a}\\right)=\\frac{8\\pi\\,G}{3}\\,[\\rmr(t)+\\rDE(t)]\\,, \\end{equation} where we always assume flat space metric. In the above expression, we have defined \\begin{equation}\\label{EffectiveDE} \\rDE(t)=\\frac{3}{8\\pi G}\\left[\\CHd\\,\\Hd(t)+\\CH H^2(t)\\,\\right]\\,. \\end{equation} It plays the role of effective DE for the generalized cosmological model. Despite the two equations above imply $\\wDE=P_{\\rm DE}/\\rDE\\equiv -1$, the entropic-force model is not  a time varying vacuum model since it does not have a $\\CC$CDM limit, that is to say, we cannot obtain a constant $\\rDE$ behavior at any time in its Hubble expansion history.  Let us emphasize at this point an issue previously mentioned in passing, but that is important when we consider the status of the basic equations (\\ref{entropica}) and (\\ref{EffectiveDE}). These equations are not necessary derived from first principles, such as a fundamental action. As a matter of fact, in the most general entropic-holographic formulations\\,\\cite{Jacobson95,Padmanabhan10} one considers that the gravitational field equations are not necessarily inferred from a fundamental action at the present macroscopic level of description. The field equations can nevertheless provide, in principle, a fully satisfactory account of all the basic phenomena known to date. In such general ``emergent gravity formulation'', the ultimate origin of gravity is claimed to lie in some fundamental degrees of freedom quite different from the metric variables, namely degrees of freedom which are completely unknown to us at present. In this sense, the field equations under discussion are also thought of as just effective field equations without having a known fundamental action behind. However, we should also remark that these particular equations pertain to a restricted formulation of the general entropic approach called ``entropic-force'' gravity, and it is our aim to discuss the phenomenological status of such restricted version of the entropic approach concerning the cosmological implications first suggested in Ref.\\cite{Easson10}. In this sense our conclusions concerning the phenomenological viability of these models cannot be extended beyond the domain of this formulation.  Implicit herein is the fact that matter is \\emph{not} conserved in the entropic-force models. Indeed, Eqs.(\\ref{entropica}) and (\\ref{GeneralizedFriedmann}) can be used to show that matter and the effective DE (\\ref{EffectiveDE}) exchange energy through the generalized conservation law: \\begin{equation}\\label{Bronstein} \\dot{\\rho}_{\\rm DE}+\\dot{\\rho}_m+3(1+\\wm)\\,H\\,\\rmr=0\\,. \\end{equation} This equation is therefore not an independent one, but it is useful to show explicitly that there is a continuous exchange of energy between matter and the generalized DE given by Eq.(\\ref{EffectiveDE}). Such equation is actually enforced by the Bianchi identity, which implies $\\nabla^{\\mu}\\tilde{T}_{\\mu\\nu}=0$ for the total energy-momentum tensor $\\tilde{T}_{\\mu\\nu}$, obtained from the ordinary matter contribution, ${T}_{\\mu\\nu}$, plus the variable DE density (\\ref{EffectiveDE}). This is in contradistinction to what it was assumed in \\cite{Easson10}. Let us note that in \\cite{Easson10b} the treatment of the covariant matter conservation law was different from the previous one by the same authors in \\cite{Easson10}, and the modified law was presented in the form \\begin{equation}\\label{correctedlaw} \\dot{\\rho}_m+3\\,(1-C_H)(1+\\wm)\\,H\\,\\rmr=0\\,. \\end{equation} However, this equation is only approximate. It can be easily derived e.g. from our equations (\\ref{GeneralizedFriedmann}) and (\\ref{EffectiveDE}) after neglecting the $\\CHd$ contribution. Below we will present the exact solution for $\\rmr$ consistent with the full conservation equation (\\ref{Bronstein}), and then (\\ref{correctedlaw}) will ensue only as a particular case.  It is not difficult to show that equations (\\ref{entropica})-(\\ref{EffectiveDE}) can be combined to produce the following differential equation for the Hubble rate: \\begin{equation}\\label{BasicdiffEquation} \\dot{H}+\\frac32\\,\\left(1+\\wm\\right)\\,\\frac{1-C_{H}}{1-\\alpha}\\,H^2 =0\\,, \\end{equation} where we have defined $\\alpha=\\frac32\\,\\left(1+\\wm\\right)\\,\\CHd$. Using the transformation from the cosmic time $t$ to the scale factor $a$ through $d/dt=aH(a)d/da$ in Eq.\\,(\\ref{BasicdiffEquation}) and defining $E(a)=H(a)/H_{0}$ we can easily present the solution of Eq.(\\ref{BasicdiffEquation}) as follows: \\begin{equation}\\label{Hentropic} E(z)=(1+z)^{3\\xim/2}\\,, \\end{equation} where $z=a(z)^{-1}-1$ is the redshift, and we have defined the important parameters \\begin{eqnarray}\\label{xiparameter} \\xim&=&\\xi(1+\\wm)\\nonumber\\\\ \\xi&=&\\frac{1-C_{H}}{1-\\alpha}\\,. \\end{eqnarray} If we consider $C_{H}=3\\CHd/2$ and $\\wm=0$, then $\\xim=\\xi=1$ and the above solution boils down to the Einstein de Sitter model.  The matter density in the entropic-force model is evaluated  from \\begin{eqnarray}\\label{eq:matterdens} \\rho_m(z)&=&-\\frac{1}{8\\pi G (1+\\wm)}\\frac{dH^2(a)}{d{\\rm ln}a}\\nonumber\\\\ &=&\\frac{1+z}{3(1+\\wm)}\\,\\rco\\,\\frac{dE^2(z)}{dz}\\,. \\end{eqnarray} Computing this expression and subsequently using it for solving Eq.\\, (\\ref{Bronstein}) we can obtain the DE density as well. The final results are: \\begin{equation}\\label{rhomrhoLentropic} \\rmr(z)=\\rco\\,\\frac{1-C_{H}}{1-\\alpha}\\,(1+z)^{3\\xim}\\equiv\\rmo\\,(1+z)^{3\\xim} \\end{equation} \\begin{equation}\\label{rhomrhoLentropic1} \\rDE(z)=\\rco\\,\\frac{C_{H}-\\alpha}{1-\\alpha}\\,(1+z)^{3\\xim} \\equiv \\rho_{{\\rm DE},0}\\,(1+z)^{3\\xim}\\,. \\end{equation}  In these expressions we identify in the matter era ($\\wm=0$, $\\xim=\\xi$) \\begin{equation}\\label{OmoOLoentropic1} \\Omo=\\frac{\\rmo}{\\rco}=\\frac{1-C_{H}}{1-\\alpha}=\\xi\\,, \\end{equation} \\begin{equation}\\label{OmoOLoentropic2} \\Omega_{{\\rm DE},0}=\\frac{\\rho_{{\\rm DE},0}}{\\rco}=\\frac{C_{H}-\\alpha}{1-\\alpha}=1-\\xi\\,, \\end{equation} which clearly satisfy $\\Omo+\\Omega_{{\\rm DE},0}=1$, as expected. Furthermore, if one neglects $\\CHd$ (hence $\\alpha\\simeq 0$) we have $3\\xim\\simeq3(1-C_H)(1+\\wm)$, and then we recover -- as a particular case-- the solution of the approximate Eq.\\,(\\ref{correctedlaw}) pointed out by the authors of Ref.\\,\\cite{Easson10b}. Finally, using Eq.(\\ref{Hentropic}) the deceleration parameter $q=-\\ddot{a}/a H^2$ can be equivalently computed from \\begin{equation}\\label{eq:qnu} q=-1-\\frac{d{\\rm ln}H}{d{\\rm ln}a}=-1+\\frac{1+z}{E(z)}\\frac{dE(z)}{dz}= -1+\\frac{3\\xim}{2} \\;. \\end{equation} It is evident that this cosmological model has no inflection point in its cosmic history, i.e. a point where deceleration can change into acceleration. This is because Eq.(\\ref{eq:qnu}) is redshift independent and hence the deceleration parameter maintains sign throughout the cosmic history. As a result the universe always accelerates or always decelerates depending on the value of $\\xim$. In the matter dominated era, namely $\\wm=0$ and $\\xim=\\xi$ , the accelerated expansion of the universe ($q<0$) is obtained for $\\xi<\\frac{2}{3}$. But in such case we would be also admitting that the universe has been accelerating forever, which is of course difficult to accept. Clearly, this feature is  the main drawback of this model.  Finally, we would like to mention once more that in this section we have discussed the background evolution via Eqs.(\\ref{entropica}), (\\ref{GeneralizedFriedmann}) of the DE models considered by the author of the entropic force models \\cite{Easson10}. Although one could define the DE in different ways, our aim was to show that using the nominal approach employed by the original authors the entropic-force model in which both kind of Hubble terms ${\\dot H}$ and $H^{2}$ appear in the effective dark energy (DE) is in trouble. In this context, the DE pressure is assumed to satisfy the vacuum equation of state, namely $P_{\\rm DE}=-\\rDE$. This is important since in the original model of \\cite{Easson10} the above equation of state was not mentioned. In our case we make this assumption, which is the minimum one we can do. In fact one can show that if another equation of state would be used the kind of problems that the entropic model has would remain. For example, if the EoS parameter for the DE would take some arbitrary constant value $\\omega_{DE}$ different from  minus one, this would just change the coefficient of $H^2$ in Eq.\\,(\\ref{BasicdiffEquation}), but the solution (\\ref{Hentropic}) would take the same form (with a slightly different expression for $\\xi_\\omega$) and hence the deceleration parameter in Eq.\\,(\\ref{eq:qnu}) would still be independent of the redshift. As a result the absence of the inflexion point between deceleration and acceleration would persist. ",
        "conclusions": "In this work we have considered the  entropic-force model\\,\\cite{Verlinde10} and reanalyzed the implications as a possible dark energy candidate in cosmology within the context of \\,\\cite{Easson10}. We have found that in its original formulation the effective dark energy implied by such model is not viable neither at the background nor at the cosmic perturbations level. Specifically, the analysis of the deceleration parameter immediately detects a fundamental problem, to wit: the expansion of the universe always accelerates or always decelerates. While this is of course a severe problem it might still be cured if such behavior were only the late-time behavior of a more complete model where the transition from deceleration to acceleration would be present. However, a detailed analysis of the cosmic perturbations furnishes also a negative result, and this may give the final blow to the model: it turns out that the existence of growing modes for structure formation is incompatible with the accelerated expansion.  In an attempt to rescue some form of the entropic-force DE model, we have also addressed a modification of its energy density by replacing the time derivative term $\\dot{H}$ with the linear term $H$ while keeping the $H^2$ one. The modified model is better behaved at the background level, for there is a transition point from deceleration to acceleration and moreover the best fit value of $\\Omo$ is near $0.3$. However, the modified model is once more unable to fit the most recent growth data on structure formation. Let us mention that if we would also suppress the $H^2$ component (i.e. if the DE density would be reduced to the linear term in the Hubble function), then the best fit value of $\\Omo$ has poor quality in the allowed region and runs completely out of range (near $1$) if we single out the CMB data.  In all these cases the absence of an additive term in the structure of the DE density lies at the root of the main difficulties with these models. This constant is essential for the transition from deceleration into acceleration. It is also essential to provide a late-time linear growth rate that can describe the structure formation data in a way comparable to the $\\CC$CDM model. If the structure of the entropic-force models is corrected with an additive term, i.e. if the DE density becomes an ``affine function'' of the Hubble terms, $\\rDE=c_0+c_1\\,H+c_{\\dot{H}}\\,\\dot{H}+c_2\\,H^2$, then these models have a well-defined $\\CC$CDM limit when the coefficients of the $H$,$\\dot{H}$ and $H^2$ terms approach zero. These modified models can be phenomenologically compatible with the cosmological data\\,\\cite{Polarski,Adria}.   Let us note that we have focused on the study of the entropic-force models of the $\\CC$-type, i.e. such that the DE density (\\ref{EffectiveDE}) has the equation of state $\\wDE=-1$, as in Ref.\\cite{Polarski}. One may consider other possible implementations\\,\\cite{Japanese12,Japanese3}, but to get a viable one it is unavoidable to introduce in all cases an additive term. Unfortunately this cannot be done in a natural way in the entropic-force formulation.  Such situation is in stark contradistinction to the running vacuum models based on the renormalization group in semiclassical gravity in curved spacetime\\,\\cite{JSP-CCReview13}. The latter types of models exist in the literature since long ago -- see \\,\\cite{RGCosmology1,RGCosmology2,Fossil07} and references therein -- and have been successfully tested\\,\\cite{Bas09c,GSBP11,PerturbVacuum1,PerturbVacuum2}. They naturally incorporate the additive term as an integration constant, namely in the ``affine'' form mentioned above, but only the even powers of $H$ (and the powers of $\\dot{H}$) are allowed, as expected from the covariant QFT formulation in a curved background.  The entropic-force models that became popular lately are actually a particular case of the running vacuum models, as clearly discussed in \\,\\cite{Polarski}. This particular case, however, proves to be phenomenologically inviable at the background and/or at the perturbations level, as we have clearly demonstrated here. Even if other formulations of the entropic models are possible the main result, namely that simple combinations of pure Hubble terms $H$,$\\dot{H}$, $H^2$ are not sufficient for a complete description of the cosmological data, stays in force."
    },
    "1402/1402.4558_arXiv.txt": {
        "abstract": "A high-energy component in the radio galaxy Centaurus A was reported after analyzing four years of Fermi data.  The spectrum of this component is described by means of a broken power law with a break energy of 4 GeV and, below and above spectral indices of $\\alpha_1$=2.74$\\pm$0.03 and $\\alpha_2$=2.09$\\pm$0.20, respectively. Also a faint $\\gamma$-ray flux at TeV energies was detected by H.E.S.S.. In this paper we show that the spectrum at GeV-TeV energies is described through synchrotron self-Compton emission up to a few GeV ($\\sim$ 4 GeV) and $\\pi^0$ decay products up to TeV energies, although the emission of synchrotron radiation by muons could contribute to the spectrum at GeV energies, if they are rapidly accelerated.  Muons and  $\\pi^0$s are generated in the interactions of accelerated protons with two populations of seed photons which were reported by Compton Gamma-Ray Observatory: one population at intermediate state emission with energy peak of 0.15 MeV and another at low state emission with energy peak of 0.59 MeV.  In addition, we show that the reported observations of ultra-high-energy cosmic rays and non high-energy neutrino detection around Centaurus A can be explained through these interactions, assuming that proton spectrum is extended up to ultra-high-energies. ",
        "introduction": "Centaurus A (Cen A) is classified as Fanaroff $\\&$ Riley Class I  \\citep{1974MNRAS.167P..31F} active galactic nucleus (AGN). At a distance of D$_z\\simeq3.8$ Mpc, it has been  one of the best studied extragalactic sources,  characterized by having an off-axis jet  of viewing angle which is estimated  as $\\sim 45^\\circ$ \\citep[see, e.g.] [and reference therein]{2006PASJ...58..211H} and two giant radio lobes of hundreds of kiloparsec.  Close to the core, this source has been imaged in radio, infrared, optical \\citep{1975ApJ...199L.139W,1976ApJ...206L..45M,1970ApJ...161L...1B,1981ApJ...244..429B},  X-ray and $\\gamma$-rays (MeV-TeV) \\citep{2003ApJ...593..169H,1999APh....11..221S,2009ApJ...695L..40A, 2010ApJ...719.1433A,2011A&A...531A..30R}. Particularly,  in the range of MeV - TeV energies,  Cen A  has been observed by Compton Gamma-Ray Observatory (CGRO) mission, Fermi Gamma-Ray Space Telescope and High Energy Stereoscopic System (H.E.S.S.) experiment.    Observations from OSSE, COMPTEL and EGRET, all of them instruments of Compton Gamma-Ray Observatory (CGRO),  imaged Cen A  from 1991 to 1995 in two states (intermediate and low).  The intermediate state exhibited its brightest peak at $0.15^{+0.03}_{-0.02}$ MeV with a flux of $4.52\\times 10^{-4}$ MeV/cm$^2$/s and luminosity of 5$\\times 10^{42}$ erg/s whereas the low one showed it at  $0.59^{+0.02}_{-0.02}$ MeV with a flux of $2.91\\times 10^{-4}$ MeV/cm$^2$/s and  luminosity of 3$\\times 10^{42}$ erg/s \\citep{1998A&A...330...97S}.  Also, for a period of 10 months, Cen A was monitored by the Large Area Telescope (LAT) on board the Fermi. The $\\gamma$-ray flux collected  was described by a power law with a photon index of 2.67$\\pm 0.10_{stat} \\pm 0.08_{sys}$ \\citep{2010ApJ...719.1433A}.  For more than 120 hr, Cen A was also observed by H.E.S.S. \\citep{2005A&A...441..465A, 2009ApJ...695L..40A}.   The spectrum was described by a power law with a spectral index of 2.7$\\pm\\, 0.5_{stat} \\pm 0.2_{sys}$ and an integral flux of $\\sim 2.6 \\times 10^{39}$ erg s$^{-1}$  above an energy threshold of $\\sim$ 250 GeV.    Recently,  \\citet{2013ApJ...770L...6S} reported  a new high-energy (HE) component in the spectrum of Cen A as a result of  analyzing four consecutive years of Fermi-LAT data.  The spectrum of this component was described using  a broken power law  with a break energy at  4 GeV and spectral indices  of $\\alpha_1$=2.74$\\pm$0.03 (below) and $\\alpha_2$=2.09$\\pm$0.20 (above). It has to be highlighted that although the spectral energy distribution (SED) has been successfully described up to GeV energy range by means of synchrotron self-Compton (SSC) emission,  there is still controversy about the emission processes that could contribute  to spectrum at GeV-TeV energies such us one-zone SSC model  \\citep{2010ApJ...719.1433A},  photo-disintegration of heavy nuclei \\citep{2013arXiv1312.7440K} and  photo-hadronic processes \\citep{2012ApJ...753...40F, 2012PhRvD..85d3012S, 2014A&A...562A..12P}. On the other hand, Pierre Auger Observatory (PAO)  associated some ultra-high-energy cosmic rays  with the direction of Centaurus A  (UHECRs)   \\citep{2007Sci...318..938P, 2008APh....29..188P} and IceCube reported 28 neutrino-induced events in a TeV - PeV energy range,  although none of them related with this direction \\citep{2013arXiv1311.5238I, 2013arXiv1304.5356I}.\\\\ In this work, we describe a hadronic model  through  the proton-photon (p$\\gamma$) interactions to describe the $\\gamma$-ray spectrum from the core emission of Cen A at the GeV - TeV energy range.   Based on the p$\\gamma$ interactions and  taking into account the seed photons at two emission states (intermediate and low), we present the synchrotron radiation by muons and $\\pi^0$ decay products.  We use the muon synchrotron emission to fit the spectrum up to $\\sim$ 126 GeV and $\\pi^0$ decay products up to a few TeV.  Finally, correlating the produced $\\gamma$-ray fluxes for both states with UHECR and neutrino fluxes, we estimate the number of expected events in PAO and IceCube, respectively. We have assumed that the proton spectrum is extended through a simple  power law up to ultra-high-energies. We hereafter use  k=$\\hbar$=c=1 in natural units and z$=0.00183 \\pm 0.00002\\simeq 0$. ",
        "conclusions": "We have proposed a leptonic and hadronic model to explain the $\\gamma$-ray spectrum at GeV - TeV energy range. In the leptonic model, we have applied the SSC emission described by \\citet{2012ApJ...753...40F} and have showed that  SSC emission only gives account of photons with energies of less than $\\approx$ 4 GeV. To explain the spectrum at higher energies than 4 GeV, we have introduced  the hadronic model,  assuming that accelerated protons in the jet interact with two photon population; at intermediate and low state emission.  Based on these p$\\gamma$ interactions, we have developed synchrotron radiation by muons and $\\pi^0$ decay products.  Considering the photon population at low state emission, we found that both Fermi and H.E.S.S. data are well described by means of photo-pion spectrum (eq. \\ref{pgammam}) with break energy of 91.3 GeV and proton luminosity of  2.2$\\times 10^{44}$ erg/s.   In this case, the contribution of  synchrotron radiation to the $\\gamma$-ray spectrum is not required, so both new Fermi and H.E.S.S. data are described by $\\pi^0$ decay products.   Additionally, we have estimated that the numbers of UHECRs and neutrinos  are  3.93 and 2.37 $\\times 10^{-2}$, respectively.  On the other hand,  assuming the photon population with intermediate state, H.E.S.S. and Fermi data are described separately;  new Fermi data with muon synchrotron radiation and H.E.S.S. data with $\\pi^0$ decay products. From this emission, we found the values of  break energy of 359.11 GeV and  proton luminosity of 2.03$\\times 10^{43}$ erg/s. Additionally, extending the proton and neutrino spectrum  again up to UHE, we estimate  that the number of UHECRs is 10.44 and that of neutrinos 0.82 $\\times 10^{-3}$.  It has to be added that although muons radiate at GeV energy range by synchrotron emission,  if muons do not radiate rapidly, this process would not be effective, then $\\pi^0$ decay products originated from low emission state would describe the whole spectrum at GeV - TeV energy range.\\\\ On the other hand, we have showed that extrapolating the proton spectrum by a simple power law up to 10$^{20}$ eV, the number of UHECRs expected is closer when considering the low (3.93) than intermediate (10.44) state emission.  Although UHECRs can hardly be accelerated up to the PAO energy range at the emission region (E$_{max}$= 40  EeV) \\citep{2010ApJ...719.1433A}, they could be accelerated during the flaring intervals and/or in the giant lobes.  It is very interesting the idea that UHECRs could be accelerated partially in the jet at energies ($<40\\times 10^{19}$ eV) and partially in the Lobes at ($E>40\\times 10^{19}$ eV).\\\\ In summary,  we have showed that hybrid leptonic SSC and hadronic processes are required to explain the $\\gamma$-fluxes at GeV- TeV energy range. We have successfully described the spectral indices, break energies and fluxes, and also the expected number of  UHECRs and neutrinos \\citep{2009ApJ...695L..40A, 2013ApJ...770L...6S,2008APh....29..188P, 2013arXiv1311.5238I}."
    },
    "1402/1402.0467_arXiv.txt": {
        "abstract": "The Alpha Magnetic Spectrometer (AMS-02) is a high energy particle detector designed to study origin and nature of cosmic rays up to a few TV from space. It was installed on the International Space Station (ISS) on May 19$^{th}$, 2011. During the first two years of operation AMS-02 performed precise measurements of the proton flux. In the low rigidity range, from 1 GV to 20 GV, the proton flux was daily measured with a statistical error less than 1\\%. In the same rigidity range a gradual decrease due to Solar modulation effect and transit variations due to Solar Flares and Coronal Mass Ejection were also observed. In the rigidity range from 20 GV up to 100 GV instead, AMS-02 data show no drastic variation and the results are consistent with other experiments. Above 100 GV, AMS-02 proton flux exhibits a single power low behavior with no fine structures nor brakes. ",
        "introduction": "Protons constitute the most abundant part of the galactic cosmic rays. Knowing their absolute flux and their spectral shape is of fundamental importance for understanding the origin and propagation of galactic cosmic rays in our Galaxy. In addition, primary protons coming from the Sun are an important probe to study solar phenomena. Proton spectral shape above 100 GeV has recently been under strong interest for the space particle community. In this paper the accurate determination of the proton flux measured by AMS-02 is presented~\\cite{bib:sada}. ",
        "conclusions": "Fig.~\\ref{pfall} shows the average proton flux over the two years of AMS-02 observation as a function of kinetic energy multiplied by the corresponding bin central value~\\cite{bib:bincen} taken to the 2.7 power and compared with previous experimental data~\\cite{bib:pamela}--\\cite{bib:crdb}. In the high energy region above 100 GeV the spectrum is consistent with a single power law spectra and shows no fine structure nor break."
    },
    "1402/1402.2979_arXiv.txt": {
        "abstract": "A new mechanism to form a magnetic pressure supported, high temperature corona above the photosphere of an accretion disk is explored using three dimensional radiation magneto-hydrodynamic (MHD) simulations. The thermal properties of the disk are calculated self-consistently by balancing radiative cooling through the surfaces of the disk with heating due to dissipation of turbulence driven by magneto-rotational instability (MRI).  As has been noted in previous work, we find the dissipation rate per unit mass increases dramatically with height above the mid-plane, in stark contrast to the $\\alpha$-disk model which assumes this quantity is a constant. Thus, we find that in simulations with a low surface density (and therefore a shallow photosphere), the fraction of energy dissipated above the photosphere is significant (about 3.4\\% in our lowest surface density model), and this fraction increases as surface density decreases.  When a significant fraction of the accretion energy is dissipated in the optically thin photosphere, the gas temperature increases substantially and a high temperature, magnetic pressure supported corona is formed. The volume-averaged temperature in the disk corona is more than $10$ times larger than at the disk mid-plane. Moreover, gas temperature in the corona is strongly anti-correlated with gas density, which implies the corona formed by MRI turbulence is patchy.  This mechanism to form an accretion disk corona may help explain the observed relation between the spectral index and luminosity from AGNs, and the soft X-ray excess from some AGNs.  It may also be relevant to spectral state changes in X-ray binaries. ",
        "introduction": "The standard $\\alpha$-disk model \\citep[][]{ShakuraSunyaev1973} is usually used to account for the thermal emission from both AGN \\citep[e.g.,][]{Krolik1999} and X-ray binaries \\citep[e.g.,][]{Doneetal2007}. However, a hard X-ray tail, or a soft X-ray component with energies much larger than that expected from thermal emission from the disk are often observed from both AGNs \\citep[][]{Elvisetal1978} and X-ray binaries \\citep[][]{RemillardMcClintock2006,Doneetal2007}. During a state transition in X-ray binaries, the hard X-ray tail may actually become stronger than the thermal emission from the accretion disk \\citep[][]{RemillardMcClintock2006}.  Hard X-rays from black hole accretion flows are usually believed to be formed via inverse Compton scattering of seed photons from the accretion disk by a hot coronae above the surface \\citep[][]{BisnovatyiBlinnikov1976,HaardtMaraschi1991,HaardtMaraschi1993,SvenssonZdziarski1994,Zdziarskietal1999}. The corona, once formed, is usually thought to be located in a compact region near the innermost stable circular orbit \\citep[e.g.][]{ReisMiller2013, DexterBlaes2013}.  However, despite the rich observational evidence for the existence of hot coronae, there is no general consensus regarding the mechanism by which such coronae are formed. Many models \\citep[][]{HaardtMaraschi1991,SvenssonZdziarski1994} simply assume that some fraction of the energy liberated by accretion is dissipated in optically thin regions.  Magnetic reconnection is often invoked as the energy dissipation mechanism in the corona \\citep[e.g.][]{Galeev1979,GoodmanUzdensky2008,UzdenskyGoodman2008}.  On the other hand, magneto-rotational instability (MRI) is now understood to be the mechanism that drives angular momentum transport and energy dissipation in black hole accretion disks \\citep[][]{BalbusHawley1991,BalbusHawley1998}.  It is therefore of great interest to explore whether the MRI naturally leads to high temperature, magnetic pressure supported corona, and if so, by what mechanism.  In fact, early numerical simulations that studied the vertical structure of accretion disks with MRI turbulence \\citep[][]{MillerStone2000} showed that indeed a strongly magnetized region was formed in the low density upper layers of the disk. However, these simulations adopted an isothermal equation of state, and therefore the gas temperature in the upper layers could only be estimated {\\em ex post facto}.  Using the flux-limited diffusion (FLD) module \\citep{TurnerStone2001} in the ZEUS code, \\cite{Hiroseetal2006} studied the vertical temperature structure of accretion disks with MRI turbulence and radiative cooling.  Although they also found magnetic pressure supported low density regions emerged in the upper layers of the disk, the gas temperature in these layers was always much smaller than at the disk mid-plane.  Detailed spectral modeling \\cite{Blaesetal2006} based on the vertical profiles of dissipation and magnetic support taken from these simulations showed a slight hardening of the spectrum compared to previous models that neglected these effects, but no sign of a much hotter corona. Subsequent models computed using the same code \\citep[][]{Kroliketal2007, Blaesetal2007, Hiroseetal2009, Blaesetal2011} that explored a wide range of ratios between radiation and gas pressure always found similar temperature profiles. At the same time, a common property of these simulations is a significant increase in dissipation rate per unit mass in the surface regions, suggesting a temperature inversion might be possible under the right circumstances.  All previously published radiation MHD simulations of the MRI have considered a large surface density so that the total electron scattering optical depth from the disk mid-plane to the surface is $\\sim 10^4$.  In this paper, we report on new radiation MHD simulations performed using the full-transport radiation transfer module in Athena \\cite[][]{Davisetal2012, Jiangetal2012}.  We find that when the disk surface density is decreased so that the total electron scattering optical depth is only $\\sim 10^2$, a strong temperature inversion forms above the disk surface due to dissipation of the turbulence in magnetically supported regions {\\em which are optically thin}. This temperature inversion is (in many respects) consistent with the observationally inferred properties of accretion disk coronae, and we will refer to it as a corona throughout this work.  Due to the thermal runaways that we generically observe in radiation pressure dominated simulations \\citep[][]{Jiangetal2013c}, we focus solely on gas pressure dominated disks in this paper.  This constraint prevents the simulations from producing the very high temperatures ($\\sim 10^9$ K) that are observationally inferred, but does not prevent an initial exploration of this potentially new mechanism for corona formation. ",
        "conclusions": "By comparing two different simulations with MRI turbulence and radiation transfer, we first confirm that when the surface density is so large that the total electron scattering optical depth is $\\sim 10^4$, there is very little dissipation in the optically thin part and the temperature in the magnetic pressure supported region is always much smaller than the disk mid-plane. When the surface density of the disk is small enough such that the total electron scattering optical depth is only $\\sim 200$, a non-negligible fraction of dissipation from MRI occurs above the photosphere, and a high temperature, magnetic pressure supported corona is self-consistently generated. Note that the disk is still in the gas pressure dominated regime and the temperature in the corona is only $\\sim 10^7$ K in this particular run. Since this is still much smaller than what is required to make the observed hard X-rays, we do not claim to have reproduced the corona of observed systems. Observationally consistent coronal models for most observed systems will require a larger fraction of the dissipation above the photosphere than found in simulation A ($\\sim 3.4\\%$). Nevertheless, we consider this a useful demonstration that MRI turbulence can produce temperature inversions in magnetically dominated surface regions in simulations with realistic thermodynamics. This also motivates future studies to see how the structures of the accretion disks will be changed if more and more energy is dissipated in the optically thin part when the disk surface density decreases further.      Figure \\ref{TPProfile} shows a notable contrast between the volume average and density-weighted average of the temperature in the surface regions.  Near the surface, there is a significant anti-correlation between the temperature and density that leads to the volume average temperature being higher than the density-weighted average.  Dissipation in this region is dominated by work done by magnetic stresses and numerical reconnection.  The associated heating is rather inhomegenous  as indicated by Figure \\ref{current} and tends to deposit a larger amount of energy per unit mass in lower density regions, giving rise to the temperature-density anti-correlation.  Such `patchy' heating is interesting because it may be necessary to explain the hard X-ray tails observed in high/soft or very high state X-ray binaries \\citep[see e.g. sec. 5.3.1 of][]{Doneetal2007}.  In these systems, the hard spectral slopes are suggestive of high temperatures (or high electron energies in a non-thermal electron distribution), but the fraction of the power emitted in the component suggests that only a small fraction of photospheric photons are scattered. Such a geometry arises naturally if a significant fraction of the heating occurs in a small fraction of the volume, as the simulations imply.  One limitation of these shearing box simulations is that they consider only a small patch of the disk, so the global structure of the accretion disk and corona cannot be studied.  Post processing of global MHD simulations with Monte Carlo radiative transfer calculations suggests that many aspects of the observed hard X-ray spectrum can be reproduced \\citep[][]{Schnittmanetal2013}.  One drawback of these calculations is that they rely on heuristic cooling function to maintain a thin disk in the simulations \\citep{Nobleetal2009}.  In the future, we hope to perform a similar analysis with global radiation MHD simulations, particularly since many of the observed systems are in the radiation pressure dominated regime.  Here, we simply discuss how these local patches of the disk would fit into standard accretion disk models.  In the $\\alpha$ disk model, the surface density of the disk decreases as the disk effective and mid-plane temperatures decrease in the gas pressure dominated regime \\citep[][]{ShakuraSunyaev1973}, consistent with both ZEUS and Athena simulation results \\citep{Hiroseetal2009b,Jiangetal2013c}.  In a global accretion disk model, this corresponds to reducing the accretion rate and total luminosity. But at the same time, we have shown that a larger fraction of energy is dissipated in the corona region. Therefore, the ratio between luminosity from high energy band and low energy band, or the spectral index, would increase when the total luminosity decreases. The transition from the high soft state to the low-hard state in X-ray binaries is a case where spectrum hardens in this manner as the luminosity decreases \\citep[][]{RemillardMcClintock2006}.  Even if the inner disk is radiation dominated at the beginning of the transition, it should become gas pressure dominated as the accretion rate drops so the surface density must decrease as well.  There is some observational evidence that during this transition, the optically thick disk is truncated at radii larger than the last stable circular orbit \\citep[][and references therein]{Doneetal2010}.  In the inner disk, the geometrically thin and optically thick disk must become a geometrically thick and optically thin ADAF-like flow \\citep{NarayanYi94}. If this interpretation is correct, there is likely to be some transition region with total electron scattering optical depth smaller than $\\sim 100$, which corresponds to the parameters adopted in our simulations.  For the radiation pressure dominated regime, the $\\alpha$ disk model predicts that the disk effective and midplane temperatures will increase as the disk surface density decreases. However, our shearing box simulations of the radiation pressure dominated regime show thermal runaways \\citep[][]{Jiangetal2013c} and it remains unclear how these runaways might saturate in a global disk. Nevertheless, \\cite{Jiangetal2013c}, found that in the radiation pressure dominated regime, the dissipation scale height changes slower than the change of density scale height. Therefore a larger fraction of energy will be dissipated at high altitude when disk surface density decreases. On one hand, this may be inconsistent with some high/soft state observations which have relatively low coronal contribution. On the other hand, the fact that the very high state (or steep power law state) has a strong tail and only seems to occur at high luminosities \\citep[][]{RemillardMcClintock2006}, suggests that the surface density of the accretion disk is only small enough to enable the formation of corona when the accretion rate is very large in the radiation pressure dominated regime. Therefore, it will still be interesting to see in the radiation pressure dominated case before the thermal runaway completely changes the structures of the disk, how corona will be formed when the surface density is decreased. This will be studied with future shearing box simulations.   Most AGN are expected to be in the radiation pressure dominated regime, except at low Eddington rates where a transition to an ADAF-like flow occurs.  Therefore, we can only speculate on the observational implications of our current, gas pressure dominated results.  We note that X-ray and optical observations of AGNs do indicate a trend of hardening of the spectral index as luminosity decreases \\citep[][]{Steffenetal2006,Justetal2007}, although the black hole mass is uncertain so it is unclear what the relevant Eddington ratios are in these systems. Recently, a new population of quasars have been found to show much weaker X-ray emission compared with other AGNs with similar luminosity range \\citep[][]{Wuetal2011,Luoetal2013}. Observationally, it is still debated whether these quasars are intrinsically X-ray weak or obscured. If the AGN X-ray emission is indeed from the corona generated by MRI turbulence as found here, then these quasars may be intrinsic X-ray weak when most of the dissipation happens inside the photosphere and the corona becomes very weak or disappears.  Finally, we note that the increased dissipation per unit mass near the disk surface may be the mechanism needed to generate the continuum soft X-ray excess inferred in some AGN \\citep{Czernyetal03,Doneetal2012}. Ultimately, testing these hypotheses will require global simulations of accretion flows in both the gas and radiation pressure dominated regimes."
    },
    "1402/1402.5597_arXiv.txt": {
        "abstract": "A statistical analysis of radial (line-of-sight) 1D-distributions of brightest cluster galaxies (BCGs) within the redshift interval $0.044 \\leq z \\leq 0.78$ and Mg II absorption-line systems ($0.37 \\leq  z \\leq 2.28$) is carried out. Power spectra and two-point radial correlation functions are calculated. It is found that both  radial distributions of spectroscopic redshifts of 52,683 BCGs and 32,840 Mg II absorption systems incorporate similar quasi-periodical components relatively to the comoving distance. Significance of the components exceeds 4$\\sigma$-level and admits an increase ($\\ga  5\\sigma$) for some broad  subsamples. For the $\\Lambda$CDM cosmological model the periodicities correspond to spatial comoving scales $(98 \\pm 3)$  and $(101 \\pm 2)$ h$^{-1}$ Mpc, respectively. These quasi-periods turn out to be close to the characteristic scale $(101 \\pm 6)$ h$^{-1}$ Mpc of the quasi-periodical component obtained earlier for the radial distribution of luminous red galaxies (LRGs). On the other hand,  the scales are close to the spatial scale corresponding to the baryon acoustic oscillations (BAOs) revealed by many authors at the last decade. Fourier transform phases obtained for the BCGs and LRGs are found to be close, while the phases calculated for the Mg II absorption systems and LRGs are opposite. Discussions of the results in a context of the BAO and large-scale structure characteristic scales are outlined. ",
        "introduction": "\\label{intro} We continue a series of works on a research of quasi-periodical scales in radial (line-of-sight) 1D-distributions of various  cosmologically remote objects. The first two papers of this series (\\citealt*{rkv07}; \\citealt{rk11}, hereafter Papers~I and ~II) investigated quasi-periodical features in the spatial-temporal distribution of 2003 and 2322 absorption-line systems (ALSs) detected in the spectra of 661 and 730 QSOs, respectively, within the cosmological redshift interval $z=0.0$--4.3. In Paper~II we revealed the existence of quasi-periodicity at a scale $\\Delta {\\rm D}_c = (108 \\pm 6)~h^{-1}$~Mpc, where D$_c$ is a line-of-sight  comoving distance to an ALS with absorption redshift $z_{abs}$. It was used the $\\Lambda$CDM cosmological model with dimensional density parameter $\\Omega_{\\rm m}=0.23$. The third  paper (\\citealt*{rkk13}, hereafter Paper~III) analyzed the radial distribution of 106,000 luminous red galaxies (LRGs) from the Sloan Digital Sky Survey (SDSS) catalogue, data release~7~(DR7), and found two most significant peaks in its power spectrum. The peaks correspond to the spatial comoving scales $(135 \\pm 12)~h^{-1}$~Mpc and $(101 \\pm 6)~h^{-1}$~Mpc at the $\\Lambda$CDM model with $\\Omega_{\\rm m}=0.25$. The latter peak is the dominant with significance $\\ga 4\\sigma$.  The appropriate scale of quasi-periodicity is in mutual accordance with the period found for the radial distribution of ALSs,  especially if one takes into account that at $\\Omega_{\\rm m}=0.25$ the spatial ALS-scale  shifts to $\\Delta {\\rm D}_c = (105 \\pm 6)~h^{-1}$~Mpc. Moreover, both scales are in agreement with a spatial scale $(102.2 \\pm 2.8)~h^{-1}$~Mpc \\citep{bl11}, which is widely accepted to be the scale of sound horizon, $r_s$, at the recombination epoch displaying itself in cosmological galactic surveys as the spatial scale corresponding to the baryon acoustic oscillations (BAOs;\\, e.g.,   \\citealt{eh98};  \\citealt*{eht98}; \\citealt{bg03}; \\citealt{eisenetal05}; \\citealt*{esw07}; \\citealt{bh09}; \\citealt{kaz10}; \\citealt{anderetal12}; \\citealt{ros12}, and references therein). To confirm statistically faintly outlined relations between the BAO-phenomenon (mainly appearing in Fourier space) and the quasi-periodical components of the radial distribution of matter it seems to be reasonable to examine different types of cosmologically remote objects in similar techniques. This is one of the causes motivating us to carry out the present statistical treatment.  In this paper we deal with the radial distributions of two independent sets of cosmological data: (a) the spectroscopic redshifts $z$ of 52,683 so-called brightest cluster galaxies (BCGs) or the most luminous galaxies among constituents of clusters \\citep{wh13} within the  interval $0.044 \\leq z \\leq 0.78$, and (b) the redshifts of 32,840 intervening Mg~II ALSs within the interval $0.37 \\leq z_{abs} \\leq 2.28$ \\citep{zm13}. The updated catalog (a) of the BCGs sampled from spectral data of the SDSS DR9 catalog \\citep{ahnetal12} is available on the website % \\footnote{http://zmtt.bao.ac.cn/galaxy\\_clusters}; the  data (b) are based on the SDSS DR7 catalog (e.g. \\citealt{abaz09}) and represented on the website% \\footnote{http://www.pha/jhu.edu/~gz323/jhusdss}. To carry out the statistical considerations we calculate the power spectra of radial distributions and separately -- two-point radial correlation functions (RCFs).   The basic value of the power spectra calculations is a radial function N(D$_c$) integrated over angles $\\alpha$ (right ascension) and $\\delta$ (declination), D$_c (z)$ is the line-of-sight comoving distance between an observer and cosmological objects under study; N(D$_c$)dD$_c$ is a number of objects inside an interval dD$_c$. The radial comoving distances are calculated in a standard way  (e.g., \\citealt*{khs97};  \\citealt{h99}) \\begin{equation} {\\rm D}_c (z_i) = {c \\over H_0}\\, \\int_0^{z_i} {1 \\over \\sqrt{{\\rm \\Omega}_m (1+z)^3 + {\\rm \\Omega}_\\Lambda}} {\\rm d}z, \\label{Dc} \\end{equation} where $i$  numerates  redshifts $z_i$ of cosmological objects in a sample, $H_0=100~h$~km~s$^{-1}$ is the present Hubble constant, $c$ is the speed of light, $c/H_0\\ =\\ 2998~h^{-1}$~Mpc; it is accepted that the dimensionless density parameters are ${\\rm \\Omega}_m = 0.25$ and ${\\rm \\Omega}_\\Lambda = 1-{\\rm \\Omega}_m = 0.75$. We use the binning approach and calculate so-called normalized radial distribution function: \\begin{equation} {\\rm NN}({\\rm D}_c) = {{\\rm N}({\\rm D}_c) - {\\rm N}_{\\rm tr}({\\rm D}_c) \\over \\sqrt{{\\rm N}_{\\rm tr}({\\rm D}_c)}}, \\label{NN} \\end{equation} where N$_{\\rm tr}$(D$_c$) is a smoothed function  filtering out the largest scales (trend function). The denominator $\\sqrt{{\\rm N}_{\\rm tr}}$ in Eq.~(\\ref{NN}) stands for the standard deviation of the Poisson statistics. Examples of the trend function are represented in Fig.~\\ref{N(D)}, where the trends are calculated by the least-squares method with using a set of parabolas, as regression functions for N(D$_c$), and independent bins $\\Delta_c = 10~h^{-1}$~Mpc.  Note that instead of the radial distribution function N(D$_c$) one can use a comoving number density $n({\\rm D}_c)$\\ =\\ N(D$_c$)/dV$_c$, where dV$_c$=$4\\pi {\\rm D}_c^2 {\\rm d D}_c$ is a comoving differential volume in the flat Universe, that is an analogue of the more conventional value $n(z)$ (e.g., \\citealt{zehavetal05};  \\citealt{kaz10}). This replacement would not change Eq.~(\\ref{NN}) and results of following calculations because in that case one should divide both numerator and denominator by the comoving volume (remind that $\\sigma(n_c)=\\sigma({\\rm N}_c)/{\\rm d V}_c$, where $\\sigma$ is the standard (mean squared) deviation).  The power spectrum is calculated for the normalized radial distribution NN(D$_c$) according to the formula (e.g.,  \\citealt{jw69}, \\citealt{sc82}) \\begin{eqnarray} \\hspace{1.0cm} {\\rm P} (k_{\\rm m}) & = & {1 \\over {\\cal N}_{\\rm b}} \\left\\{ \\left[ \\sum_{j=1}^{{\\cal N}_{\\rm b}}  {\\rm NN}_j \\cos(k_{\\rm m} {\\rm D}_{c, j}) \\right]^2   \\right. \\nonumber                           \\\\ & +  & \\left.    \\left[ \\sum_{j=1}^{{\\cal N}_{\\rm b}} {\\rm NN}_j \\sin (k_{\\rm m} {\\rm D}_{c, j}) \\right]^2 \\right\\}, \\label{Pk} \\end{eqnarray} where ${\\cal N}_{\\rm b}$ is a number of bins, $j=1,2,...{\\cal N}_{\\rm b}$ is a numeration of bins, ${\\rm D}_{c, j}$ is a location of bin centers, $k_{\\rm m}=2\\pi {\\rm m}/{\\rm D}_c^{\\rm L}$ is a wave number, corresponding to an integer harmonic number m$=1,2,...$,  ${\\rm D}_c^{\\rm L}$ is the whole comoving interval. Note that Eq.~(\\ref{Pk}) represents just a square of the  discrete  Fourier transform, $F_{\\rm m}=F(k_{\\rm m})$, responsible for the m-th harmonic (e.g., \\citealt{g75}): $|{\\rm F}_{\\rm m}|^2 = {\\rm Re}^2({\\rm F}_{\\rm m}) + {\\rm Im}^2({\\rm F}_{\\rm m}))$.  We determine also the two-point RCF, $\\xi(\\delta {\\rm D}_c)$, for the samples under consideration in an unconventional way (e.g., Papers~I and II). The variable $\\delta {\\rm D}_c$ is a radial comoving distance between components of any pair of the objects in the samples. Note that members of different pairs with fixed $\\delta {\\rm D}_c$ may correspond to different mutual comoving distances. Using the estimator by \\citet{ls93} one can introduce \\begin{equation} \\xi(\\delta {\\rm D}_c) = {{\\rm DD} \\over {\\rm RR}} - 2 {{\\rm DR} \\over {\\rm RR}} + 1, \\label{Xi} \\end{equation} where DD = DD$(\\delta {\\rm D}_c)$ is a number of observed pairs of cosmological objects separated by the radial comoving distance $\\delta {\\rm D}_c$ belonging to a range $\\delta {\\rm D}_c \\pm \\tilde{\\Delta}_c/2$, where $\\tilde{\\Delta}_c$ is a chosen accumulative bin width; DR = DR($\\delta {\\rm D}_c$) is a number of cross pairs between points of the real sample and a random sample simulated for the same intervals of D$_c$ and  the same smoothed function (trend) as the real one; RR = RR($\\delta {\\rm D}_c$) is a number of pairs counted up only for the random sample.  We concentrate on the consideration of periodical components  incorporated in the radial distributions of the samples relatively to the D$_c$-variable. We treat these components as {\\it quasi}-periodicities because of limited intervals of line-of-sight distances used for both types of the objects, as well as a dependence of peak positions and amplitudes on sample variations. In Sections~2 and 3 we present the results of statistical analysis of galaxy clusters, traced by the BCGs and absorption systems Mg~II, respectively. In Section~4  we confront Fourier transform phases of the main quasi-periodical components with the phase of the dominant quasi-periodical component of the radial distribution of LRGs (Paper~III). Conclusions and discussion of the results in a context of the large-scale structure (LSS) of the matter distribution in the Universe are given in Section~5. ",
        "conclusions": "\\label{concl} The main results of the statistical treatment of  52,683 brightest cluster galaxies (BCGs) from the SDSS DR9 within the redshift interval $0.044  \\leq  z  \\leq  0.78$ \\citep{wh13} and 32,840 Mg~II absorption-line systems (ALSs) from SDSS DR7 within the interval $0.37 \\leq z \\leq 2.28$ \\citep{zm13} can be summarized as follows:  (1)  The radial (line-of-sight) 1D-distribution of  BCGs incorporates  a significant quasi-periodical component with respect to the smoother radial selection function (trend). The dominant peak in the power spectra demonstrates the  periodicity corresponding to a spatial characteristic scale $\\Delta {\\rm D}_c = (98 \\pm 3)~h^{-1}$~Mpc at a level of significance $\\ga 4\\sigma$. Still more prominent spectral peak,  $\\ga 5\\sigma$, corresponding to the same scale is obtained in the power spectrum calculated for a subsample of the 41,420~most luminous BCGs with absolute magnitudes $M_{\\rm r} \\leq -23.01$. Let us emphasize, however, that the mean value of  relative amplitudes of the oscillations averaged over the whole interval D$_c^{\\rm L}$ for the latter subsample turns out to be small  $\\sim 0.04$.  (2) The existence of the quasi-periodicity in the radial distribution of BCGs has been confirmed also by the two-point  radial (1D) correlation function $\\xi (\\delta {\\rm D}_c)$ displaying the periodicity in the explicit form.   (3) The radial distribution of Mg~II~ALSs also includes a significant quasi-periodical component with respect to the trend. The dominant peak in the power spectra corresponds to a scale $\\Delta {\\rm D}_c = (101 \\pm 2)~h^{-1}$~Mpc at the same significance level $\\ga 4\\sigma$ as for the BCGs. We also obtain an enhancement of the peak height for the radial distribution of  26,772~Mg~II~systems separated out by two partly compatible conditions:  W$_r(2796) < 0.75\\AA$  or $\\delta z \\leq 0.0001$ at W$_r(2796)  \\geq 0.75 \\AA$, where $\\delta z$ is an accuracy of redshift $z_{abs}$ detections. In the latter case the mean value of  relative amplitudes of the oscillations averaged over the whole interval  D$_c^{\\rm L}$ is also small  $\\sim 0.06$. Note that similar scale $\\sim 100~h^{-1}$~Mpc of the C~IV~ALSs correlation function within a range $1.2 < z  < 4.5$ was found also by \\citet*{qvby96}.  (4) The two-point  radial (1D) correlation function $\\xi (\\delta {\\rm D}_c)$ also visualizes the quasi-periodical component in the radial distribution of Mg~II~ALSs corresponding to the same scale as the peak in the power spectra.  (5) Both the quasi-periodical scales found here in the radial distributions of two different types of cosmological objects are consistent with the quasi-periodicity scale $(101 \\pm 6)~h^{-1}$~Mpc found earlier (Paper~III) in the radial distribution of the luminous red galaxies (LRGs). The phases of the main periodical components, calculated as reciprocal Fourier transforms, turn out to be approximately close for the BCGs and LRGs and approximately opposite (shifted at a phase angle $\\sim \\pi$) for the Mg~II~ALSs and LRGs.  (6) All considered quasi-periodical scales (partly overlapping) are consistent with the scale $(102.2 \\pm 2.8)~h^{-1}$~Mpc of the baryon acoustic peak in the two-point spatial (3D-monopole) correlation function of the LRGs \\citep{bl11}. It is likely that there are implicit relations between the quasi-periodicities found here in the radial distributions of cosmological objects and the characteristic scale of sound horizon, $r_s$, at the recombination epoch (so called  {\\it standard ruler}; e.g., \\citealt{bg03}; \\citealt{percetal07}; \\citealt{esw07}; \\citealt{percetal10}; \\citealt{anderetal12}). In that case further confirmation and specification of the quasi-periodicity in radial distributions would  become a simple and alternative tool for an express-analysis of vast arrays of  cosmological data. On the other hand, the proximity of the scales of the baryon acoustic peak and those treated in the present paper may evidence in favour of an assumption that primordial acoustic perturbations responsible for the BAO-scale could carry traces of a partly ordering formed at some early epochs, e.g., recombination or radiation-matter equipartition.  For instance, one can imagine a weak  tendency of galaxies to be distributed as a continuous structure of higher and lower densities (a set of cells) formed by standing acoustic waves (e.g.,  \\citealt{zn83},  \\citealt*{hss97}, \\citealt{esw07}) with dominant characteristic scale $\\sim r_s$. This structure may display itself  in some apparent arrangement of the matter distribution, as a consequence of a possible relatively dense packing of the cells. In such a context the hypothesis of a weak (linear) partial ordering of matter imprinted in the epoch of recombination (at $z \\la z_{rec} \\sim 10^3$) may relate to a conception of a quasi-regular network formed by voids and galaxy superclusters at the nonlinear stage ($z \\ll z_{rec}$), which was discussed, e.g., in a series of papers by Einasto et al. (e.g.,  \\citealt{einetal97a, einetal97b, einetal97c}; \\citealt{ein00}; \\citealt{tagoetal02}; \\citealt{einetal11}).  An illustration of the ordering possibilities was represented also in Paper~II. There was simulated a model of partly ordered structure of points (galaxies or galaxy clusters) forming a cloud-like scattering around the vertices of a simple cubic (SC) lattice. For this model it was produced a mean power spectrum averaged over 100 random realizations of the radial distribution function. The mean spectrum displays a significant peak  at $k=2\\pi/l_{SC}$, where $l_{SC}$ is a lattice constant  and resembles the  spectrum shown in Fig.~\\ref{Pk_Xi_BCG}(a). That reflects a general property of the radial distribution of the lattice points with an arbitrary centre of the space. In this case the resultant radial distribution exhibits a periodical behaviour with a constant amplitude of oscillations (e.g., \\citealt{km63}).  To elucidate the relation between the oscillations observable in 3D-power spectrum, displaying the BAO phenomenon,  and quasi-periodical component in the radial distribution, revealed in this work, we carried out some additional simulations. To be specific, we produced a square matrix of complex Fourier amplitudes $F(k_x, k_y)$ in a 2-dimensional $k$-space, where $k_x$ and $k_y$ are two projections of the wave vector ${\\bf k}$, $k^2=k_x^2+k_y^2$, and Re$(F(k_x, k_y))$ and Im$(F(k_x, k_y))$ obey the Gaussian distributions, which are modulated by harmonic variations with smoothly decreasing (damped) amplitudes. The damping is similar to the observable baryon oscillations in 3D-$k$-space (e.g., \\citealt{eh98};  \\citealt{bg03}; \\citealt{anderetal12}). Then we transformed the generated in such a way square matrix of $F(k_x, k_y)$ into the distribution of points in the real 2D-space and applied the analysis described in Introduction to random realizations of the distribution obtained.  Our preliminary results have shown that at some decreasing dependences of the amplitudes at $k=k_1, k_2, k_3 ...$, where $k_1, k_2=2 k_1, k_3=3 k_1 ...$ -- positions of the oscillation maxima, the power spectrum  Eq.~(\\ref{Pk}) of the radial distribution displayed a set of descending narrow peaks at the same $k=k_1, k_2, k_3 ...$. Note that subsequent  peaks cab be less significant than the first ones. Moreover, the radial distribution turns out to be sensitive to harmonic loops of a 2D-$k$-distribution transforming  them to narrower peaks in the related 1D-$k$-space, i.e., intensifying one (or a few) discrete radial harmonic(s). On the other hand, in the real 2D-space we obtain only one noticeable peak of the standard (multi-central) correlation function. That reminds the well-known peak of the 3D-correlation function (e.g.,  \\citealt{eisenetal05};  \\citealt{kaz10}; \\citealt{bl11};  \\citealt{ros12}) specific for the BAO phenomenon in the real space. Being confirmed these results could reconcile  the baryon oscillations in the $k$-space with quasi-periodicities of the radial distribution in the real space.  It is likely as well that the quasi-periodicities formulated in items (1)-(5) concern  with the well known results by \\citet{beks90} widened and confirmed by \\citet{szetal93}. Their pencil-beam surveys near both the Galactic poles within a redshift interval $z \\la 0.5$ displayed a periodicity at a scale 128~$h^{-1}$~Mpc. Quasi-periodicities in two-dimensional patterns obtained on the base of the Las Campanas Redshift Survey data had been also found by \\citet{landyetal96} as a  power-spectrum peak corresponding to a scale of $\\sim 100~h^{-1}$~Mpc in the spatial distribution of galaxies. This scale is rather close to the scales found here in the radial distributions of the  BCGs and Mg~II ALSs, and found earlier for the LRGs. Possible relations of the cited periodicities with the spatial scale of BAOs $\\sim 100~h^{-1}$~Mpc were discussed by \\citet{ehss98} and \\citet{sz98}.\\ \\citet{szetal93} emphasized the special role of rich-statistics pencil-beam surveys in  bringing out the phase correlations (regularities) of matter density fluctuations. They expected  systematic decrease and broadening of the power spectrum peaks for wider angle regions of observational samples.  Note, however, that the region of the SDSS DR9 survey on the sky (e.g., \\citealt{ahnetal12}) includes only the north galactic pole. Additional elimination from the sample of all objects within solid angles at $\\theta \\leq 5^\\circ$ around the pole axis does not effect noticeably the radial quasi-periodical component of the rest sample of BCGs. Thus the effect of total radial distribution considered here in a sense differs from the effects discussed by \\citet{beks90} and \\citet{szetal93}. Although, it is likely that there is some implicit relation between these two approaches which still have to be manifested in a future.  In reality quasi-periodical components of the matter distribution may appear as a consequence of two (or a few) factors, e.g.,  effects of clumping of galaxies and their clusters (e.g.,  \\citealt{kp91}) and/or quasi-ordering of the matter spatial distribution (discussed above). It is likely that the modern surveys of the sky deal with a complex mixture of both (or more) the factors.  In any case, the quasi-periodicity revealed here being confirmed probably traces some spatial partly-ordered structure imprinted in the LSS of matter in the early Universe.  \\textit"
    },
    "1402/1402.5418_arXiv.txt": {
        "abstract": "During the stalled-shock phase of our three-dimensional, hydrodynamical core-collapse simulations with energy-dependent, three-flavor neutrino transport, the lepton-number flux ($\\nu_e$ minus $\\bar\\nu_e$) emerges predominantly in one hemisphere. This novel, spherical-symmetry breaking neutrino-hydrodynamical instability is termed LESA for ``Lepton-number Emission Self-sustained Asymmetry.'' While the individual $\\nu_e$ and $\\bar\\nu_e$ fluxes show a pronounced dipole pattern, the heavy-flavor neutrino fluxes and the overall luminosity are almost spherically symmetric. Initially, LESA seems to develop stochastically from convective fluctuations. It exists for hundreds of milliseconds or more and persists during violent shock sloshing associated with the standing accretion shock instability. The $\\nu_e$ minus $\\bar\\nu_e$ flux asymmetry originates predominantly below the neutrinosphere in a region of pronounced proto-neutron star (PNS) convection, which is stronger in the hemisphere of enhanced lepton-number flux. On this side of the PNS, the mass-accretion rate of lepton-rich matter is larger, amplifying the lepton-emission asymmetry, because the spherical stellar infall deflects on a dipolar deformation of the stalled shock. The increased shock radius in the hemisphere of less mass accretion and minimal lepton-number flux ($\\bar\\nu_e$ flux maximum) is sustained by stronger convection on this side, which is boosted by stronger neutrino heating due to \\hbox{$\\langle\\epsilon_{\\bar\\nu_e}\\rangle > \\langle\\epsilon_{\\nu_e}\\rangle$}. Asymmetric heating thus supports the global deformation despite extremely nonstationary convective overturn behind the shock. While these different elements of the LESA phenomenon form a consistent picture, a full understanding remains elusive at present. There may be important implications for neutrino-flavor oscillations, the neutron-to-proton ratio in the neutrino-heated supernova ejecta, and neutron-star kicks, which remain to be explored. ",
        "introduction": "Nonradial hydrodynamic instabilities play an important role in the postbounce dynamics of collapsing stellar cores on their way to the onset of supernova (SN) explosions. They also accompany the formation and cooling of the proto-neutron star (PNS). These phenomena include convection in the PNS, large-scale convective overturn below the stalled shock wave during the accretion-shock phase, and the standing accretion shock instability (SASI).  We presently add a new phenomenon to this list which we call LESA for ``Lepton-number Emission Self-sustained Asymmetry.'' Its most conspicuous manifestation is lepton-number flux ($\\nu_e$ minus $\\bar\\nu_e$) emission primarily in one hemisphere, but it also involves dipolar asymmetry of convection inside the PNS and different strengths of large-scale convective overturn below the stalled shock in both hemispheres.  We first recall these traditional spherical-symmetry breaking instabilities and begin with convection. Prompt postshock convection leads to the decay of the negative entropy and electron-number gradients left behind by the weakening bounce shock and the shock-breakout burst of electron neutrinos, respectively \\citep{Burrows_1992,Janka_1993,Mueller_1994}. It fosters shock expansion and acts as a source of gravitational-wave emission for a period of some ten milliseconds after core bounce \\citep{Mueller_1997,Mueller_2013}. Inside the nascent neutron star (NS), i.e.\\ below (and possibly also around) the neutrinosphere, Ledoux convection was expected to occur because of the negative lepton number gradient produced by the inward progression of the deleptonization wave associated with neutrino losses \\citep{Epstein_1979,Burrows_1988}. In two-dimensional (2D) simulations (axial symmetry), PNS convection was first studied by \\citet{Keil_1996} and later again by \\citet{buras_06b} and \\citet{Dessart_2006}.  \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=\\textwidth]{f1.pdf} \\caption{Lepton-number flux ($\\nu_e$ minus $\\bar\\nu_e$) for our $11.2\\,M_\\odot$ model as a function of direction for the indicated times post bounce. The latitudes and longitudes, indicated by dotted lines, correspond to the angular coordinates of the polar grid of the numerical simulation. The flux in each panel is normalized to its average, i.e., the quantity   $(F_{\\nu_e}-F_{\\bar\\nu_e})/\\langle F_{\\nu_e}-F_{\\bar\\nu_e}\\rangle$ is color coded. The lepton-number emission asymmetry is a large-scale feature, which at later times has clear dipole character. The black dots indicate the positive dipole direction of the flux distribution, the black crosses mark the negative dipole direction. The dipole track between 70 and 340\\,ms is shown as a dark-gray line.  Once the dipole is strongly developed, its direction remains essentially stable and shows no correlation with the $x$-, $y$-, and $z$-axes of the numerical grid. The dipole direction is also independent of polar hot spots, which are persistent, local features of moderate amplitude and an artifact connected with numerical peculiarities near the z-axis as coordinate singularity of the polar grid.} \\label{fig:leptonskyplots} \\end{center} \\end{figure*}  Large-scale convective overturn below the stalled shock was predicted by \\citet{bethe_90} and confirmed by the first 2D hydrodynamical simulations \\citep{Herant_1992,Herant_1994,Burrows_1995,Miller_1993,Janka_1995,Janka_1996,Mezzacappa_1998} as well as in 3D \\citep{Fryer_2002,Fryer_2004}. The driving force is a negative entropy gradient that develops a few tens of milliseconds after core bounce in the neutrino-heating region between the gain radius (where neutrino heating begins to exceed neutrino cooling) and the stalled shock. Convective flows stretch the dwelling time of matter in the gain layer and thus increase the energy deposition by neutrinos. This effect can provide crucial support to the delayed neutrino-heating mechanism: multi-D simulations can yield explosions even when spherically symmetric models fail (e.g., \\citealt{Janka_1996,Murphy_2008,Nordhaus_2010,Hanke_2012,Dolence_2013,Couch_2013,Couch_2013a}).  The delayed-explosion mechanism is also aided by SASI, which leads to violent shock sloshing motions. This effect expands the shock, increases the gain layer and, again, can enhance the efficiency of neutrino-energy deposition \\citep{marek_09} even when convection is weak or its growth is suppressed because of a small shock-stagnation radius and correspondingly fast infall velocities in the gain layer \\citep{Foglizzo_2006,Scheck_2008}. This nonradial instability was first observed in 2D simulations with a full 180$^\\circ$ grid \\citep{Janka_1996,Mezzacappa_1998,Janka_2003,Janka_2004}, but not immediately recognized as a new effect beyond large-scale convection. It was unambiguously identified in 2D hydrodynamical simulations of idealized, adiabatic (and thus non-convective) postshock accretion flows \\citep{Blondin_2003}. SASI was found to possess the highest growth rates for the lowest-order (dipole and quadrupole) spherical harmonics \\citep{Blondin_2006,Foglizzo_2007,Iwakami_2008} and to give rise to spiral-mode mass motions in 3D simulations \\citep{Blondin_2007,Iwakami_2009,Fernandez_2010,Hanke_2013} or in 2D setups without the constraint of axisymmetry \\citep{Blondin_2007,Yamasaki_2008,Foglizzo_2012}. The instability can be explained by an advective-acoustic cycle of amplifying entropy and vorticity perturbations in the cavity between accretion shock and PNS surface \\citep{Foglizzo_2002,Foglizzo_2007,Scheck_2008,Guilet_2012} and has important consequences for NS kicks \\citep{Scheck_2004,Scheck_2006,Nordhaus_2010b,Nordhaus_2012,Wongwathanarat_2010,Wongwathanarat_2013} and spins \\citep{Blondin_2007,Rantsiou_2011,Guilet_2013}, quasi-periodic neutrino emission modulations \\citep{Marek_2009,Lund_2010,Tamborra_2013}, and SN gravitational-wave signals \\citep{Marek_2009,Murphy_2009,Mueller_2013}.  Here we report the discovery of a new type of low-mode nonradial instability, LESA, which we have observed in 3D hydrodynamical simulations with detailed, energy-dependent, three-flavor neutrino transport using the \\textsc{Prometheus-Vertex} code.  Our current portfolio of simulated 3D models includes an $11.2\\,M_\\odot$ model that shows violent large-scale convection, but no obvious signs of SASI activity during the simulated period of postbounce evolution, and two models (20$\\,M_\\odot$ and 27\\,$M_\\odot$) in which episodes of SASI alternate with phases of dominant large-scale convection \\citep{Hanke_2013,Tamborra_2013,Tamborra_2014}. While all models exhibit LESA, with different orientations of the emission dipole, the clearest case is the $11.2\\,M_\\odot$ model, because the new effect is not overlaid with SASI activity.  To provide a first impression of our new and intriguing phenomenon we show in Fig.~\\ref{fig:leptonskyplots} the distribution of lepton-number emission ($\\nu_e$ minus $\\bar\\nu_e$) for the $11.2\\,M_\\odot$ model over the stellar surface at postbounce (p.b.) times of 148, 169, 210, and 240\\,ms. In each panel, the lepton-number flux is normalized to the instantaneous average and the color scale covers the range from $-$0.5 to 2.5 of this relative measure. We indicate the positive dipole direction with a black dot, the negative direction with a cross. We also show the track of the positive dipole direction as a dark-gray line, ranging from 70\\,ms p.b., where the dipole begins forming, to the end of the simulation at 340\\,ms. While at 148\\,ms the dipole pattern is not yet strong---a quadrupole component is clearly visible and the dipole is still building up as we will see later---the subsequent snapshots reveal a strong dipole pattern with large amplitude: In the negative-dipole direction, the lepton-number flux is around zero, and even negative in some small regions, whereas in the positive direction it is roughly twice the average and even larger in some small regions.  LESA is a large and conspicuous effect concerning the deleptonization flux. At the same time, it is associated with more subtle dipole deformations of other quantities. In particular, the dipolar lepton-number emission is linked to anisotropic PNS convection, which leads to an aspherical electron distribution in the PNS mantle layer. The emission dipole is additionally fed by a hemispheric mass-accretion asymmetry, which might cause the one-sided enhancement of PNS convection. This accretion asymmetry in turn is a consequence of a dipolar shock deformation that deflects the accretion flow preferentially to one hemisphere.  Despite vigorous and highly time-variable, nonstationary convective overturn stirring the postshock region, the shock deformation and mass-accretion asymmetry are maintained for hundreds of milliseconds by a neutrino-heating asymmetry that is tightly linked to the neutrino-emission anisotropy: Because $\\bar\\nu_e$ have somewhat harder spectra than $\\nu_e$, neutrino heating on the side of a relatively higher $\\bar\\nu_e$ flux (the side with lowest lepton-number flux and lowest mass accretion rate) is stronger and sustains the dipolar shock deformation that produces the hemispheric asymmetry of the postshock accretion flow.  This preliminary interpretation suggests that LESA is not a purely hydrodynamical phenomenon, in contrast to the traditional instabilities, but depends on a complex interaction of hydrodynamic mass flow and neutrino emission and heating.  Our new effect is predominantly a hemispheric asymmetry in these quantities and as such is not a generic 3D effect, but it has not been previously reported in the context of 2D simulations. Actually, indications of LESA may be present in the O-Ne-Mg core explosions of \\citet{Wanajo_2011} and in a $15\\,M_\\odot$ explosion model of \\citet{mueller_12}. However, it seems difficult to make a strong case for this neutrino-hydrodynamical instability based on 2D simulations. The constraint of axisymmetry resticts any dipole asymmetry to the direction of the polar axis of the grid. This symmetry axis, where reflecting boundary conditions are imposed, defines a preferred direction and has various problematic consequences. It tends to artificially create hemispheric differences by deflecting the converging flows either inward or outward, and the grid axis also directs shock-sloshing motions. Strong bipolar motions of the postshock accretion layer in most 2D simulations, where shock expansion and contraction alternate violently between the poles, could interfere with the LESA phenomenon or even create asymmetries of different nature.  In the following, we first describe briefly, in Sect.~\\ref{sec:numerics}, the numerical setup of our 3D simulations and the overall properties of our three progenitor models. In Sect.~\\ref{sec:results} we study various manifestations of our new phenomenon, ranging from dipole deformations of neutrino-emission properties to asymmetric PNS convection. Next we turn in Sect.~\\ref{sec:feedback-loop} to more subtle manifestations in the form of asymmetric accretion and neutrino heating, which however form a feedback loop and as such are the driving engine of the overall effect. In Sect.~\\ref{sec:explanation} we string the different elements together and provide an overall scenario that involves the outer feedback mechanism consisting of asymmetric mass accretion and neutrino heating and the inner mechanism of asymmetric electron-density distribution and PNS convection. We conclude in Sect.~\\ref{sec:conclusions} with a summary and a discussion of possible implications.   \\begin{figure*}[t!] \\begin{center} \\hbox to\\textwidth{\\includegraphics[width=.33\\textwidth]{f2a.pdf}\\hfil \\includegraphics[width=.33\\textwidth]{f2b.pdf}\\hfil \\includegraphics[width=.33\\textwidth]{f2c.pdf}} \\vskip1.5pt \\hbox to\\textwidth{\\includegraphics[width=.33\\textwidth]{f2d.pdf}\\hfil \\includegraphics[width=.33\\textwidth]{f2e.pdf}\\hfil \\includegraphics[width=.33\\textwidth]{f2f.pdf}} \\caption{Snapshots of the convective overturn activity during the accretion phase of the 11.2\\,$M_\\odot$ model at the indicated p.b.\\ times. Visualized are surfaces of constant entropy: The bluish, semi-transparent envelope is the SN shock, the red-orange-yellow surfaces are entropy structures in the postshock region. Neutrino-heated high-entropy matter  expands in typical mushroom-like, buoyant plumes. These features are  highly time dependent with bubbles continuously appearing, expanding and  rising, and disappearing to give way to a new generation of convective  plumes. The direction of strongest lepton-number emission, i.e., the largest excess of the radiation of $\\nu_e$ compared to $\\bar\\nu_e$ (see text), points  towards the observer, above the $x$-$y$-plane at roughly 45$^\\circ$  inclination to the $z$-axis. A corresponding global, persistent dipolar deformation of the accretion shock is present but can hardly be recognized without a detailed analysis.} \\label{fig:3dconv11} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have described a new type of nonradial deformation mode that we discovered in our 3D stellar core-collapse simulations using energy-dependent, three-flavor neutrino transport, applied to progenitor stars of 11.2, 20, and 27\\,$M_\\odot$. During the first 100--150\\,ms of postbounce evolution, a long-lasting, only slowly evolving dipolar neutrino-emission asymmetry establishes itself. It persists throughout the  postbounce accretion phase of the stalled SN shock for hundreds of ms, even in those models that show violent SASI activity for parts of their postbounce evolution. This multi-dimensional phenomenon has not been identified unambiguously in any previous 2D simulation. We call the new phenomenon LESA for ``Lepton-number Emission Self-sustained Asymmetry.'' LESA emerges from an instability, i.e., the spherically symmetric state is not stable and the asymmetry grows from any perturbation, no matter how small. In contrast to convection or SASI, the nature of LESA is not simply hydrodynamical, but rather a neutrino-hydrodynamical instability and as such the first of its kind identified in the SN context.  The dipole mode manifests itself most conspicuously in the lepton-number flux ($\\nu_e$ minus $\\bar\\nu_e$), whose dipole amplitude can reach 100\\% of its $4\\pi$ directional average, i.e., in one direction the lepton-number flux can exceed twice the average, in the opposite direction it can be even somewhat negative ($\\bar\\nu_e$ number flux exceeds that of $\\nu_e$). While dipole amplitudes of the individual $\\nu_e$ and $\\bar\\nu_e$ number and energy fluxes of 10--20\\% can be observed in our 3D calculations of the 11.2\\,$M_\\odot$ progenitor (and somewhat smaller ones in the 27\\,$M_\\odot$ model), heavy-lepton neutrinos as well as the sum of $\\nu_e$ and $\\bar\\nu_e$ exhibit only percent-level dipolar luminosity asymmetries.  The neutrino-emission dipole originates from a hemispheric asymmetry of the electron distribution in the mantle layer of the PNS interior to the neutrinospheres, although the density stratification is perfectly spherical and concentric in these regions of extreme gravitational field strength. In the hemisphere of higher $\\nu_e$ flux, convection in the deeper layers of the PNS mantle is enhanced compared to the opposite side and dredges up electrons more efficiently from the dense, lepton-rich central core. These electrons are mixed outward towards the neutrinospheres and lead to less deleptonized conditions in one hemisphere, supporting the higher $\\nu_e$ flux. The convective activity seems to be instigated and fostered by shear flows, turbulent motions, and gravity waves caused by the violent impacts of accretion streams in the neutrinospheric region. This connection is suggested by the observed correlation between PNS-convection asymmetry and a considerable asymmetry of the mass-accretion flow, which is stronger in the hemisphere of enhanced PNS convection. The accretion asymmetry also contributes to the lepton-number emission dipole, because the accretion streams carry electron-rich matter and thus supply the PNS with fresh lepton number, but this effect is subdominant.  The global accretion asymmetry is maintained by anisotropic neutrino heating in the gain layer behind the stalled SN shock, because $\\bar\\nu_e$ leave the neutrinosphere with higher mean energy than $\\nu_e$. Therefore, neutrino heating is stronger on the side of lower lepton-number flux, despite the nearly isotropic energy flux of $\\nu_e$ plus $\\bar\\nu_e$. Stronger neutrino heating enhances convective overturn in the postshock layer, pushes the shock to a larger stagnation radius and thus produces a dipolar deformation of the shock surface. This shock deformation in turn deflects the accretion flow falling through the shock and, in the time-averaged sense, amplifies the accretion flux to the hemisphere of the PNS facing away from the greater shock radius (Fig.~\\ref{fig:cartoon}).  Anisotropic neutrino heating therefore establishes a feedback mechanism between the neutrino-emission asymmetry on one side and shock deformation and accretion asymmetry on the other. It thus mediates a complex, mutual dependence between lepton-number transport by neutrino fluxes and convection inside the PNS on the one hand and anisotropic convective overturn in the gain layer on the other. This feedback, which involves neutrinos as crucial players, allows the global dipolar asymmetry to become a self-sustained phenomenon, which exists in stable conditions over many dynamical timescales despite the presence of vigorous and highly time-dependent convective overturn in the postshock region and even through phases of violent SASI activity. Stochastic fluctuations of this convective overturn or of the convection in the PNS mantle are probably responsible for initiating the development of the hemispheric asymmetry. The convective SN core seems to be generically unstable against such a dipolar mode of asymmetry. LESA and SASI are independent phenomena, but the latter can influence the former in a complex manner, depending on the relative orientations of the associated vector directions.  LESA could have important implications for a variety of physical processes in the SN core, most importantly nucleosynthesis in the neutrino-heated ejecta, and potentially NS~kicks and neutrino-flavor conversion.  Concerning nucleosynthesis, we recall that charged-current reactions of $\\nu_e$ and $\\bar\\nu_e$ with neutrons and protons do not only heat the SN blast wave but also set the n/p ratio in the neutrino-driven outflow and thus determine the nucleosynthesis conditions in the innermost SN ejecta. Since $\\nu_e$ absorption converts neutrons to protons while $\\bar\\nu_e$ captured on protons create neutrons, the exposure to a higher flux of $\\bar\\nu_e$ tends to neutronize the expelled matter. The lepton-number emission asymmetry could persist until the explosion sets in or even beyond, fuelled by continued anisotropic PNS convection and/or asymmetric accretion beyond the onset of the explosion. In this case a considerable hemispheric asymmetry of the electron fraction in the ejecta could be expected with possibly favorable conditions for neutron-rich material in the direction where the $\\bar\\nu_e$ flux has its emission maximum. We speculate that the O-Ne-Mg core explosion of \\citet{Wanajo_2011}, where convective overturn but not SASI played a role, may be a case where we have encountered consequences of the dipolar lepton-emission asymmetry in a 2D model. The moderate hemispheric differences of $Y_e$ in the early neutrino-driven wind (with variations in the range $0.40\\lesssim Y_e \\lesssim 0.54$) in this simulation might give an impression of the corresponding effects that could be obtained in future 3D explosions.  Asymmetric neutrino emission imparts a recoil on the nascent NS. We assume a dipolar emission anisotropy of the total neutrino-energy loss rate (the summed contributions of $\\nu$ and $\\bar\\nu$ of all flavors), \\begin{equation} \\frac{{\\mathrm{d}}\\dot E_\\mathrm{tot}}{{\\mathrm{d}}\\Omega} = \\frac{\\dot E_\\mathrm{tot}}{4\\pi} \\left ( 1 + \\alpha\\,\\cos\\vartheta \\right ) \\,, \\label{eq:lumdipole} \\end{equation} where $\\alpha$ is the dipole amplitude and $\\vartheta$ the zenith angle relative to the dipole direction. In this case the NS acceleration is \\begin{equation} a_\\mathrm{NS} = \\frac{1}{M_\\mathrm{NS}c}\\int_{4\\pi}\\mathrm{d}\\Omega\\, \\frac{{\\mathrm{d}}\\dot E_\\mathrm{tot}}{{\\mathrm{d}}\\Omega} \\,\\cos\\vartheta = \\frac{\\alpha}{3\\,c}\\,\\frac{\\dot E_\\mathrm{tot}}{M_\\mathrm{NS}} \\,. \\label{eq:nsacceleration1} \\end{equation} With $\\dot E_{53}\\equiv\\dot E_\\mathrm{tot}/(10^{53}\\,\\mathrm{erg\\,s}^{-1})$ and $M_{1.5} \\equiv M_\\mathrm{NS}/(1.5\\,M_\\odot$) we obtain \\begin{equation} a_\\mathrm{NS} \\approx 37\\,\\,\\frac{\\alpha}{0.01}\\, \\frac{\\dot E_{53}}{M_{1.5}}\\,\\,\\mathrm{km\\,s}^{-2}\\,. \\label{eq:nsacceleration2} \\end{equation} If $\\alpha\\sim 0.01$--0.02 and the duration of the emission asymmetry lasts only some hundred milliseconds, corresponding to the duration of the accretion phase, the recoil velocity will not exceed several 10\\,km\\,s$^{-1}$, depending on the time-integrated neutrino-energy release $E_\\mathrm{tot}$. But the NS kick velocity could reach 100--200\\,km\\,s$^{-1}$ for a canonical value of $3\\times 10^{53}\\,$erg for the NS gravitational binding energy, if the neutrino-emission dipole continues to exist for the whole period of PNS neutrino cooling. With the luminosity maximum coinciding with the $\\bar\\nu_e$ emission peak in our 3D simulations, the NS acceleration will point in the direction of the strongest $\\nu_e$ emission. Even in the optimistic (and highly speculative) case that the neutrino-emission dipole survives for seconds, however, the estimated recoil velocity is dwarfed by those that can typically be expected from the ``gravitational tug-boat mechanism'' associated with the anisotropic ejection of matter in 3D simulations of SN explosions \\citep{Wongwathanarat_2010,Wongwathanarat_2013}.  Another small but amusing mechanical consequence of LESA is an angular momentum transfer, i.e., a spin-up of the nascent NS. Weak interactions violate parity maximally, implying that a relativistic $\\nu_e$ has negative helicity and carries the spin angular momentum $-\\hbar/2$ relative to its direction of motion, whereas a $\\bar\\nu_e$ has positive helicity and carries $+\\hbar/2$. We denote the lepton-number flux dipole amplitude with $A_{\\rm Dipole}$, i.e., the quantity plotted in Fig.~\\ref{fig:dipole11}. The angular momentum transfer rate then has the magnitude \\begin{equation} \\dot J=\\frac{\\hbar}{2}\\,\\frac{A_{\\rm Dipole}}{3}\\,, \\end{equation} because our normalization of the dipole amplitude implies that it is three times the total lepton-number flux projected on the dipole direction. With a typical value $A_{\\rm Dipole}=3\\times10^{56}\\,{\\rm s}^{-1}$ and recalling that $\\hbar=1.054\\times10^{-27}\\,{\\rm cm}^2\\,{\\rm g}\\,{\\rm s}^{-1}$ we find a typical angular-momentum transfer rate of $\\dot J=5\\times10^{28}\\,{\\rm cm}^2\\,{\\rm g}\\,{\\rm s}^{-2}$. We recall that the moment of inertia of a homogeneous sphere with mass $M$ and radius $R$ is $I=2 M R^2/5$ and its angular momentum is $J=I\\,2\\pi f$ with $f$ the rotation frequency. Using as benchmark values $M=1.5\\,M_\\odot$ and $R=15\\,{\\rm km}$, a typical PNS spin-up rate is of the order of $\\dot f\\sim 3\\times10^{-18}\\,{\\rm Hz}\\,{\\rm s}^{-1}$, a very small value indeed. A much larger effect, however, is non-radial neutrino emission that can transfer orbital angular momentum and cause a spin-down of a rotating NS as first shown by \\citet{Mikaelian_1977} and \\citet{Epstein_1978}.  Our entire study has ignored neutrino flavor conversion. The LESA phenomenon depends on a subtle hemispheric asymmetry of neutrino heating rates. Certainly this effect would be modified if the $\\nu_e$ and $\\bar\\nu_e$ fluxes would partially swap flavor with $\\nu_x$ and $\\bar\\nu_x$ which have different spectral properties and different number fluxes. Moreover, flavor conversion also modifies the neutrino influence on the n/p ratio in the context of nucleosynthesis as first pointed out by \\citet{Qian_1993} and \\citet{Qian_1995}.  The thorniest problem in the context of neutrino flavor conversion is the role of neutrino-neutrino refraction, which causes many complications because of the feedback of flavor conversion on itself \\citep{Duan_2010}. The justification for ignoring flavor conversion in the dense region below the stalled shock wave is the so-called multi-angle matter effect, which tends to suppress self-induced flavor conversion \\citep{Esteban_2008,Sarikas_2012,Raffelt_2013,Chakraborty_2011,Saviano_2012,Chakraborty_2014}. In particular, \\citet{Dasgupta_2012} have studied the onset of self-induced flavor conversion in a somewhat simplified 3D model. These authors conclude that flavor conversion always begins outside the shock wave.  It is conceivable, however, that these conclusions must be modified in the LESA context, notably in those directions where the lepton number flux is small. The asymmetry between the $\\nu_e$ and $\\bar\\nu_e$ number fluxes, sometimes denoted with the parameter $\\epsilon$, is a crucial ingredient for the stability of the neutrino flux in flavor space~\\citep{Esteban_2007}. Moreover, the LESA phenomenon also modifies the electron-density profile which defines the matter effect for neutrino flavor oscillations. Therefore, it remains to be verified that flavor conversion is indeed irrelevant for the LESA phenomenon.  The observations reported in this paper raise many interesting questions. It is obvious that the LESA phenomenon needs much more work and analysis to develop a full understanding, especially concerning how asymmetric PNS convection is stimulated by asymmetric mass accretion and how SASI and LESA interact. Many of our explanations remain tentative and require further confirmation.  In particular, further studies are necessary to reveal how generic the lepton-number emission dipole is. How exactly does its amplification work, what is the underlying mechanism? What is the role of the PNS contraction in this context, and, if it is important, how sensitive is the phenomenon to the nuclear equation of state and the neutrino opacities? How big is the saturation amplitude of LESA and what does it depend on? Could its amplitude be larger than in our present 3D simulations and could its dipolar neutrino-heating asymmetry affect the onset of neutrino-driven SN explosions? How long does LESA last? Does LESA require the inner engine of hemispherically asymmetric PNS convection as a necessary ingredient, or is the role of PNS convection only that of an amplifier of the lepton-number flux asymmetry and that of a stabilizing factor which enables long-term memory? A feedback loop seems possible that is solely based on the outer engine, in which the asymmetric accretion and lepton-number emission on the one hand is intertwined with the dipolar neutrino-heating asymmetry and shock deformation on the other. Answers to these questions are indispensable to reach firm conclusions on the importance of LESA for NS kicks, the explosion mechanism and asymmetries, and SN nucleosynthesis.  Linear growth studies would help to develop deeper insights. It would be especially useful to construct toy models that capture the essential parts of the mechanism but reduce its complexity and thus allow a better control of the interplay of its different components. Such studies would be particularly useful because numerical models are always prone to artifacts. For example, it is unclear whether our findings are affected by approximations involved in our treatment, e.g., the ray-by-ray-plus transport, which does not include nonradial neutrino fluxes, the use of a monopole description of the gravitational potential, or the spherical core of 10\\,km radius, which fixes the PNS to its location at the grid center. It is therefore essential that other groups investigate their neutrino-hydrodynamics simulations for hints of effects similar to our LESA phenomenon.   \\bigskip\\noindent {\\bf Note added.} After our work had been circulated in preprint form, the authors of \\citet{Couch_2013a} informed us that they also found LESA features in their 3D models. On the other hand, \\citet{Dolence_2014} do not find any evidence for LESA in their 2D simulations (see their Fig.~11, where they plot the evolution of the dipole-to-monopole ratio for the lepton-number flux). Performing a similar analysis for our own 2D models, we do find a strong LESA effect, commensurate with our 3D findings. However, in 2D we see an oscillatory, high-amplitude lepton-flux asymmetry, which can change its north-south orientation on timescales of order 10\\,ms. This behavior differs from typical features of LESA in 3D, like its persistence and directional stability, and asymmetric accretion and heating are difficult to diagnose in such a strongly time-variable situation. While a more detailed future study of our 2D models may shed more light on this question, we conclude that the appearance or not of LESA is not directly related to the dimensionality of the simulation.  One difference between our works is that we treat neutrino transport in the ray-by-ray-plus approximation, an approach also used by \\citet{Couch_2013a}, whereas \\citet{Dolence_2014} use a multi-dimensional flux-limited diffusion scheme, leading them to speculate that LESA is an artifact of the ray-by-ray technique. We note, however, that there are numerous other differences, including neutrino interaction rates, and that our 2D models have rather different physical characteristics at comparable epochs, in particular differ strongly in the evolution of the shock radius and actually explode. Understanding the differences in these results requires a systematic study of the influence of all physics inputs and their numerical representation. The ray-by-ray approximation ignores lateral radiative neutrino transport, although it includes, of course, advective transport of neutrinos trapped in the stellar medium (see \\citealt{buras_06a}). It seems rather implausible that lateral radiative diffusion, caused by relatively weak transverse gradients in the PNS, could suppress a global dipole asymmetry that would otherwise occur. Similarly, it is difficult to imagine that the accretion and lepton-emission asymmetries on the largest possible (hemispheric) scale could be wiped out by angular smoothing of the neutrino heating in the gain layer on much smaller scales (cf.\\ \\citealt{Sumiyoshi_2014}). Such findings would be intriguing in their own right. In any case, a satisfactory answer to this question depends on a true physical understanding of the LESA phenomenon and on observing it in simulations where the neutrino treatment does not have the shortcomings of the ray-by-ray-plus approximation."
    },
    "1402/1402.0717_arXiv.txt": {
        "abstract": "{Short millimeter observations of radio-loud active galactic nuclei (AGN) offer an excellent opportunity to study the physics of their synchrotron-emitting relativistic jets, from where the bulk of radio and millimeter emission is radiated. On one hand, AGN jets and their emission cores are significantly less affected by Faraday rotation and depolarization than at longer wavelengths. On the other hand, the millimeter emission of AGN is dominated by the compact innermost regions in the jets, where the jet can not be seen at longer wavelengths due to synchrotron opacity.} {We present the first dual frequency simultaneous 86\\,GHz and 229\\,GHz polarimetric survey of all four Stokes parameters of a large sample of 211 radio loud active galactic nuclei, designed to be flux limited at 1\\,Jy at 86\\,GHz.} {The observations were most of them made in mid August 2010 using the XPOL polarimeter on the IRAM 30\\,m millimeter radio telescope.} {Linear polarization detections above $3\\sigma$ median level of $\\sim1.0$\\,\\% are reported for 183 sources at 86\\,GHz, and for 23 sources at 229\\,GHz, where the median $3\\sigma$ level is $\\sim6.0$\\,\\%. We show a clear excess of the linear polarization degree detected at 229\\,GHz with regard to that at 86\\,GHz by a factor of $\\sim1.6$, thus implying a progressively better ordered magnetic field for blazar jet regions located progressively upstream in the jet. We show that the linear polarization angle, both at 86 and 229\\,GHz, and the jet structural position angle for both quasars and BL~Lacs do not show a clear preference to align in either parallel or perpendicular directions. Our variability study with regard to the 86\\,GHz data from our previous survey points out a large degree variation of total flux and linear polarization in time scales of years by median factors of $\\sim1.5$ in total flux, and $\\sim1.7$  in linear polarization degree --maximum variations by factors up to $6.3$, and $\\sim5$, respectively--, with 86\\,\\% of sources showing linear polarization angles evenly distributed with regard to our previous measurements.} {} ",
        "introduction": "\\label{Intr}  Among all classes of active galactic nuclei (AGN), radio-loud AGN are characterized by having pairs of relativistic jets of highly energized, magnetized plasma that are ejected along the rotational poles of the SMBH--disk system \\citep[e.g.,][]{Blandford:1982p6283}. In particular, blazars, the most exotic class of radio-loud AGN, stand out by wild variability of their non-thermal jet-emission from radio to $\\gamma$-rays. Members of this class include BL~Lacs, and flat spectrum radio quasars (FSRQ, the high power version of the former). The remarkable properties of blazars are affected by strong relativistic effects that beam their radiation and shorten the variability time scales. These properties include superluminal motions \\cite[e.g.,][]{Gomez:2001p201, Jorstad:2005p264, Agudo:2012p17361}, Doppler boosted emission of the jet pointing at a small angle to the observer \\citep[e.g.,][]{Kadler:2004p5629}, substantial changes in flux and linear polarization, sometimes correlated on several spectral ranges, on time scales from months \\citep[e.g.,][]{Bach:2006p354,Agudo:2011p14707,Agudo:2011p15946} to hours \\citep[e.g.,][]{Ackermann:2010p13506, Tavecchio:2010p14858}, as well as changes of jet structure \\cite[e.g.,][]{2007AJ.134.799J, Agudo:2007p132}, and polarized synchrotron and inverse-Compton emission all along the spectrum \\cite[e.g.,][]{Marscher:2010p11374,Jorstad:2010p11830,Agudo:2011p14707,Agudo:2011p15946,Wehrle:2012p17706}. AGN with powerful jets oriented at larger angles to the line of sight (i.e. radio galaxies, the remaining prominent class of radio-loud AGN) suffer less from such relativistic effects, therefore displaying smaller apparent luminosities and longer time scales of variability, although they are thought to be the same astrophysical objects as BL~Lacs and FSRQ, the latter being the better oriented counterparts of the former \\citep{Urry:1995p6301}.  Explaining the existence and physical properties of jets from AGN remains one of the greatest current challenges in high-energy astrophysics. Observations of radio-loud AGN at short millimeter wavelengths offer an excellent opportunity with regard to radio centimeter observations to study these objects for several reasons. First, like AGN radio emission, millimeter emission is well known to be radiated predominantly from the jet, with essentially no contribution from the host AGN or its surroundings, which facilitates the interpretation of the measurements in terms of jet physics. Second, AGN jets, and in particular their millimeter emitting cores, are significantly less affected by Faraday rotation and depolarization than at longer wavelengths \\citep[see e.g.][]{Zavala:2004p138,Agudo:2010p12104}. Third, the millimeter emission of AGN is dominated by the compact innermost regions in the jets \\citep[e.g.][]{2007AJ.134.799J,Lee:2008p301}, where the jet can not be seen at longer wavelengths due to synchrotron opacity.  In 2005, we used the XPOL polarimeter \\citep{2008PASP..120..777T} on the IRAM 30\\,m Telescope to perform the first 3.5\\,mm (86\\,GHz) polarization survey of radio loud AGN over a large (145 source) sample \\citep{Agudo:2010p12104}. Within the available data, we detected linear polarization above $3\\sigma$ levels of $\\sim1.5$\\,\\% for 76\\,\\% of the sample. Our results pointed out an excess, by a factor of $\\sim2$, of 86\\,GHz linear polarization degree with regard to that at radio wavelengths (15\\,GHz), suggesting that either the region of bulk millimeter emission has a better ordered magnetic field, or that the radio emission is strongly Faraday depolarized with regard to that at millimeter wavelengths. We also reported a trend of decreasing luminosity towards larger linear polarization degrees within our entire sample, perhaps indicating lower magnetic field order for larger luminosity jets. Moreover, in contrast to what it was found at radio wavelengths \\citep[see][ and references therein]{Lister:2005p261}, we do not find a relation between the linear polarization angle and the jet structural position angle in either quasars of BL~Lacs. Confirmation of this result would imply that no AGN class matches the conditions to show linear polarization angles either parallel or perpendicular to the jet, i.e. that they are intrinsically non-axisymmetric.  Aiming to improve the statistical confidence with regard to previous studies, we performed a new simultaneous 3.5 (86\\,GHz) and 1.3\\,mm (229\\,GHz) AGN polarimetric survey with the XPOL polarimeter on the IRAM 30\\,m Telescope in polarimetric mode over an updated and increased sample of 211 sources that was designed to be 1\\,Jy total-flux limited. Our new observations were also contemporaneous with current large surveys at other spectral ranges, therefore allowing for multi spectral range studies. The dual frequency configuration of our observing program also allowed us to search for large Faraday rotation effects in bright AGN, whereas the comparison with our previous survey allowed us to study the total flux and polarization AGN variability at short millimeter wavelengths in bright AGN. In this paper we describe this new survey and the results from the analysis of the resulting data. The comparison of the data from this new survey with $\\gamma$-ray measurements from the Fermi-LAT 2-year AGN Catalog \\citep{Ackermann:2011p17255}, as well as a detailed statistical study of their correlation are presented in a separate paper (Agudo et al., A\\&A, in prep). ",
        "conclusions": "We have presented the first simultaneous 3.5 and 1.3\\,mm polarization survey of radio loud AGN on a large sample of 211 sources, dominated by blazars.  Our total flux measurements show that almost all measured sources (96\\,\\%) have spectral index $\\tilde\\alpha_{86{\\rm,}229}^{\\rm{Q}}<0$ --most of them with optically thin spectrum-- between 86 and 229\\,GHz with median spectral indexes  $\\tilde\\alpha_{86{\\rm,}229}^{\\rm{Q}}=-0.75$ for quasars and $\\tilde\\alpha_{86{\\rm,}229}^{\\rm{B}}=-0.56$ for BL~Lacs. In contrast, the 15 to 86\\,GHz spectral index is distributed towards flatter and more mildly optically-thin spectral-indices both for quasars and BL~Lacs.  We detect linear polarization above $3\\sigma$ levels for 88\\,\\% and 13\\,\\% of the sources detected at 86 and 229\\,GHz, respectively. We find the distributions of ${m}_{\\rm{L}}$ for quasars and BL~Lacs to be significantly different at both 86 and 229\\,GHz with median linear polarization degrees $\\tilde{m}_{\\rm{L,86}}^{\\rm{Q}}=3.2$\\,\\% and $\\tilde{m}_{\\rm{L,229}}^{\\rm{Q}}=7.7$\\,\\% for quasars, and $\\tilde{m}_{\\rm{L,86}}^{\\rm{B}}=4.6$\\,\\% and $\\tilde{m}_{\\rm{L,229}}^{\\rm{B}}=10.2$\\,\\% for BL~Lacs. The differences between quasars and BL~Lacs is explained by the fact that quasars have, in general, smaller viewing angles than BL Lacs, which produce stronger depolarization in quasars if the magnetic fields in the jets of blazars are not homogeneously distributed along the jet axis or the jets themselves are not axisymmetric.  We show for the first time that the 229\\,GHz linear polarization degree is, in general, also a factor $\\sim1.6$ larger than the one at 86\\,GHz. We also confirm that the 86\\,GHz is in general stronger than that at 15\\,GHz by a factor of $\\sim1.6$. This evidence implies that the magnetic field is progressively better ordered in blazar jet regions located progressively upstream in the jet.  An anti-correlation between the millimeter luminosity and the linear polarization degree is confirmed here for the entire source sample and quasars at 86\\,GHz. We attribute this relation to purely relativistic and orientation effects, i.e. sources whose jets are better oriented to the line of sight display larger luminosities (because of their larger Doppler boosting) and also lower linear polarization degrees (because of cancellation of orthogonal polarization components along the line of sight).  Unlike theoretical predictions from axially symmetric jet models, we show here an essentially inexistent relation between $\\chi_{86}$ and the jet structural position angle for both quasars and BL~Lacs  to distribute at a preferential misalignment angle \\citep[see also][]{Agudo:2010p12104}. Only a small 17\\,\\% of quasars tend to have $0^{\\circ}\\le|\\chi_{86}-\\phi_{\\rm{jet}}|\\lesssim30^{\\circ}$. This result seems to be reproduced also at 229\\,GHz, although at this frequency the small number of sources detected in polarization do not allow to obtain robust conclusions. These results imply that either the magnetic field or the emitting particle distributions (or both) responsible for the synchrotron radiation in the innermost regions where the short mm emission is radiated in blazars have a markedly non axisymmetric character in general.  Comparison of our data with the 86\\,GHz measurements presented in \\citet{Agudo:2010p12104} shows a considerably high total flux variability factor ($:S_{86}^{\\rm{var}}$) with median variations $\\bar{:S_{86}^{\\rm{var}}}\\sim1.5$ in time scales of years, with a 19\\,\\% of sources showing $:S_{86}^{\\rm{var}}\\gtrsim2$, and even one source (\\object{0422+004}) experiencing  extreme total flux variations by $:S_{86}^{\\rm{var}}=6.3$. The 86\\,GHz linear polarization degree is also equally highly variable, with median variability factor $\\bar{:m_{L,86}^{\\rm{var}}}\\sim1.7$, and 34\\,\\% of sources displaying $:m_{L,86}^{\\rm{var}}\\gtrsim2$ with maximum variations by a factor of $\\sim5$. The 86\\,GHz linear polarization angle is also highly variable in time scales of years. Except for a small fraction (14\\,\\% of sources) of sources that tend to conserve the same polarization angle, most sources show drastically different polarization angle --evenly distributed among all possible angles between $0^{\\circ}$ and $90^{\\circ}$--  in such time scales."
    },
    "1402/1402.3302_arXiv.txt": {
        "abstract": "We search for the lensing signal of massive filaments between 135,000 pairs of Luminous Red Galaxies (LRGs) from the Sloan Digital Sky Survey. We develop a new estimator that cleanly removes the much larger shear signal of the neighboring LRG halos, relying only on the assumption of spherical symmetry. We consider two models: a ``thick'' filament model constructed from ray-tracing simulations for $\\Lambda$CDM model, and a ``thin'' filament model which models the filament by a string of halos along the line connecting the two LRGs. We show that the filament lensing signal is in nice agreement with the thick simulation filament, while strongly disfavoring  the thin model. The magnitude of the lensing shear due to the filament is below $10^{-4}$. Employing the likelihood ratio test, we find a 4.5$\\sigma$ significance for the detection of the filament lensing signal, corresponding to a null hypothesis fluctuation probability of $3 \\times 10^{-6}$. We also carried out several null tests to verify that the residual shear signal from neighboring LRGs and other shear systematics are minimized. ",
        "introduction": "One of the most striking features of $N$-body simulations for $\\Lambda$CDM structure formation scenario is the network of filaments into which dark matter particles arrange themselves. Some attempts to quantify this network have been made \\citep{s2011, cwj2014}. Other work has attempted to study the largest filaments, those between close pairs of large dark matter halos \\citep{ckc2005}. Such filaments are likely the easiest to identify in data, e.g., \\citet{zdm2013} look for overdensities in the galaxy distribution between close pairs of galaxy clusters. However, since filaments include both dark and luminous matter, weak lensing techniques are  useful to understand the entire structure: \\citet{dwc2012} and \\citet{jjk2012} both identify single filaments by focusing on a weak lensing analysis of individual cluster pairs.  In this study we measure the weak lensing signal of filaments between stacked Luminous Red Galaxy (LRG) pairs in Sloan Digital Sky Survey (SDSS) data. The mass distribution and therefore weak lensing shear in the neighborhood of LRG pairs is dominated by the massive halos themselves.  Methods which aim at filament detection, e.g., \\citet{mm2013}, may have a degeneracy with the signal from these nearby halos. In the face of this degeneracy, we construct an estimator of the lensing signal which removes the shear due to these halos, assuming only that they are spherically symmetric. We will show that this technique is sufficient to obtain a detection, and some physical implications on filament size and shape can be extracted by comparison to filament models. Systematic errors which are expected to be spherically symmetric with respect to the halos, such as intrinsic alignments, are nulled simultaneously.  Other work has attempted to estimate the feasibility of weak lensing stacked filament detection. \\citet{mm2013} make optimistic choices for survey parameters and find that $\\sim 2 - 4\\sigma$ detections are possible for single clusters but state that their method has difficulties in application to stacked filament detection. In another study \\citet{mkm2010} use lens and source redshifts that make their lensing strength a factor of 2 greater than ours, and a galaxy number density at least a factor of 30 higher. The lower mass limit of their stacked clusters is $M_{200} = 4 \\times 10^{14} M_\\odot / h$, much larger than the dark matter halos associated with our LRGs. With these parameters, they estimate that $\\sim 20$ cluster pairs are necessary to obtain a detection.  Filaments can also be characterized using the language of higher-order correlations. In this case, one would describe the filament as the part of the matter-matter-matter three point function in the neighborhood of the halos forming a cluster pair. A detection of the halo-halo-matter 3-point function around such cluster pairs was made using the Red Cluster Survey \\citep{sps2008}. More recently \\citet{simonetal:13} used CFHTLens survey to measure the galaxy-galaxy-shear correlation function and attempted to measure the average mass distribution around galaxies. This measurement was done by subtracting off the two point contribution of the lensing signal. As these authors discovered, the three-point signal peaks at the cluster locations. However, for our purposes of identifying filaments, such a location of the three point function's peak makes the technique of two-point subtraction unsatisfactory. Just as our nulling estimator removes two-point contributions which are spherically symmetric about the halo centers, it also removes any three-point contribution which is centered on these points.  We use a different approach to these authors. We use data from the SDSS which is shallower, but covers 8,000 square degrees. This allows us to use $\\sim$ 100,000 pairs of LRGs and obtain a significant detection with a new estimator of filament lensing.  The outline of this paper is as follows: section \\ref{sect:method} describes the basic nulling technique for removing spherically symmetric components. In section \\ref{sect:data} the LRG pair catalog and background source shear catalog used in this work are described. Section \\ref{sect:theory} describes our mock LRG catalog and ray-tracing simulations, as well as an alternative thin-filament model. In section \\ref{sect:results} we present our main results, including the filament measurement from data and simulations, null tests, and the likelihood ratio statistic. Finally, section \\ref{sect:discussion} discusses the implications of our results as well as directions for future work.  Throughout this work we use cosmological parameters $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $\\sigma_8 = 0.83$. ",
        "conclusions": "\\label{sect:discussion}  We have presented a technique for the statistical measurement of properties of the dark matter filaments linking close pairs of galaxies and clusters. Our method uses an empirical approach which relies on few assumptions to cancel out the contribution of spherical halos in the data. Applying this technique to pairs of LRGs in SDSS we detect residual shear patterns of magnitude $\\sim10^{-4}$, about one order of magnitude smaller than the LRG halo signal at 5 Mpc/h from the halo center, and two orders of magnitude smaller than the LRG signal at its virial radius. Employing a likelihood ratio test, we attribute this signal to filamentary structures with a detection significance higher than $4\\sigma$. The small signal, the possible systematics in SDSS imaging, and the dominance of the LRG halos make the measurement especially challenging. We have carried out several null tests which provide evidence that our nulling technique is robust and that systematic errors are minimized. Applying the same LRG pair selection and lensing estimator to simulations, we find an excellent match with the data. With a width of a few $\\mpch$ the stacked filament shape is determined to be thicker than the LRG halos at its end points.  There are several approximations and sources of error in our analysis. \\begin{itemize} \\item The stacking of hundreds of thousands of LRG pairs leads to a smearing of the mass distribution. This means that we cannot make definitive statements about the typical (individual) filament structures in the universe, in particular the limits we obtain on the thickness of the filament only apply to the stacked profile. \\item The binning scheme we use to null out the contribution of spherical halos also nulls {\\it part} of the signal from the filaments we seek to measure. In this sense our estimator is sub-optimal (though it would be challenging to construct an optimal estimator while subtracting the LRG halo signal). \\item The calibration of the shear, which relies on a correction for the smearing due to the PSF, introduces a redshift dependent bias that propagates to the filament mass estimate. Uncertainties in the photometric redshifts of background galaxies have a similar effect. Both are below the 10 percent level and  smaller than our statistical uncertainties. \\item Redshift space distortions: the line of sight separation of the LRGs is uncertain owing to their relative peculiar velocity. We have attempted to account for it by adding RSD to the simulations (see Sec.~\\ref{sect:mocks}). \\item The inevitable contamination of the LRG sample with other galaxies and stars leads to a dilution of the signal. This should be controlled to better than the 10\\% level. \\end{itemize}  In future work several improvements can be made that address the above points. In addition, forward modeling of the measurement can be done using simulations and the halo model, so that comparisons can be made without use of our nulling technique. Such an approach may allow for more detailed tests of the halo model and of filamentary properties, though care will need to be exercised to distinguish systematic errors. Finally, an interesting complement to our study is to compare the mass distribution inferred from lensing shears with the distribution of foreground galaxies and hot gas."
    },
    "1402/1402.4522_arXiv.txt": {
        "abstract": "How would negative energy density affect a classic Friedmann cosmology?  Although never measured and possibly unphysical, certain realizations of quantum field theories leaves the door open for such a possibility. In this paper we analyze the evolution of a universe comprising varying amounts of negative energy forms. Negative energy components have negative normalized energy densities, $\\Omega < 0$. They include negative phantom energy with an equation of state parameter $w<-1$, negative cosmological constant: $w=-1$, negative domain walls: $w=-2/3$, negative cosmic strings: $w=-1/3$, negative mass: $w=0$, negative radiation: $w=1/3$ and negative ultralight: $w > 1/3$. Assuming that such energy forms generate pressure like perfect fluids, the attractive or repulsive nature of negative energy components are reviewed. The Friedmann equation is satisfied only when negative energy forms are coupled to a greater magnitude of positive energy forms or positive curvature. We show that the solutions exhibit cyclic evolution with bounces and turnovers.The future and fate of such universes in terms of curvature, temperature, acceleration, and energy density are reviewed. The end states are dubbed ``big crunch,\" `` big void,\" or ``big rip\" and further qualified as ``warped\",``curved\", or ``flat\",``hot\" versus ``cold\", ``accelerating\" versus ``decelerating\" versus ``coasting\". A universe that ends by contracting to zero energy density is termed \"big poof.\" Which contracting universes ``bounce\" in expansion and which expanding universes ``turnover\" into contraction are also reviewed. ",
        "introduction": "The classic Friedmann equation of energy  is typically written in the form \\cite{weinberg72, peebles88,peacock99,mukhanov05} \\begin{equation} \\label{FriedmannOld} H^2 = {8 \\pi G \\over 3} \\rho - {k\\over R^2} , \\end{equation} where $H$ is the Hubble parameter, $G$ is the gravitational constant, $\\rho$ is the energy density, $R$ is a scale factor of the universe, and $k$ is a dimensionless constant related to the curvature of the universe. The speed of light $c$ is set to unity. The Hubble parameter $H = {\\dot R} / R = {\\dot a / a}$, where $a$ is the dimensionless scale factor of the universe such that $a = R / R_0$ and $R_0$ is the scale factor of the universe at some canonical time $t_0$.   In \\cite{nemiroff08} (hereafter Paper I), the average energy density $\\rho$ was explicitly expanded into all possible stable {\\it positive} energy forms, including phantom energy, cosmological constant, domain walls, cosmic strings, compact matter, radiation, and hypothetical energy forms collectively dubbed ultralight. These ``stable\" energy forms are considered different than energy fields such as scalar fields in that each has a well defined local form and evolves in the universe as an constant power of $a$. All energy forms evolve as an integer power of the scale factor $a$.   In this paper, we consider stable energy forms with {\\it negative} energy density, $\\rho < 0$. To the best of our knowledge, some of these energy forms have never been discussed explicitly. It is not being suggested that any of these energy forms are presently important in the universe, but their  theoretical inclusion presents cosmologically interesting scenarios. A universe comprising mostly of negative energy density component will perhaps violate various energy conditions of general relativity, and may even violate the second law of thermodynamics~\\cite{hawking}. But some extensions of quantum field theory and recent studies of Cassimir force experiments indicate  that energy conditions in relativity need only be satisfied on a global scale, or on an average measure, leaving the possibility of violation of the very same energy conditions locally or for a small period of time~\\cite{epstein65,roman,ford96,pfenning,helfer,visser00, graham,fewster}. Therefore, for  example, if some physical process that might have lasted only for a very short time compared to the age of the physical universe and may also have violated relativistic energy conditions could still play a role in the evolution of the universe by virtue of producing small amounts of negative energy components in the past that could become relevant at a later time in the evolution of such universes.  Following Paper I, the total energy density $\\rho$ can be expanded into its component stable forms such that \\begin{equation} \\label{rho} \\rho = \\sum_{n=-\\infty}^{\\infty} \\rho_n^+ \\, a^{-n} + \\sum_{m=-\\infty}^{\\infty} \\rho_m^- \\, a^{-m}, \\end{equation} where $\\rho_n^+$ is the usual positive energy density typically considered and $\\rho_m^-$ is the negative cosmological energy density. In general, in this work, the subscript $n$ and the (sometimes redundant superscript plus sign) will refer to a positive energy density, while the subscript $m$ (and superscript minus sign) will refer to a negative energy density. Note that the values of $\\rho_n^+$ and $\\rho_m^-$ remain fixed to their values at $a = 1$ and do not change as the universe evolves---hence their designation as ``stable\".   As usual, a normalizing (positive) critical density $\\rho_c$ is defined such that $\\rho_c = 3 H^2/(8 \\pi G)$ so that $\\Omega = \\rho/\\rho_c$. Dividing both sides of Eq.~(\\ref{rho}) by $\\rho_c$ at $a=1$, we obtain \\begin{equation} \\label{OmegaSum} \\Omega  = \\sum_{n=-\\infty}^{\\infty} \\Omega_n^+ \\, a^{-n} +  \\sum_{m=-\\infty}^{\\infty} \\Omega_m^- \\, a^{-m} . \\end{equation} Following above convention, also used in Paper I, $\\Omega$ depends on $a$ and hence the time $t$. Note that $\\Omega$ is different from $\\Omega_{total}$, the sum of all of the stable $\\Omega_n^+$ and $\\Omega_m^-$ values, where \\begin{equation} \\Omega_{total} = \\sum_{n=-\\infty}^{\\infty} \\Omega_n^+  + \\sum_{m=-\\infty}^{\\infty} \\Omega_m^-  . \\end{equation}   The curvature term in Eq. (\\ref{FriedmannOld}) can be written in terms of more familiar quantities. We divide Eq. (\\ref{FriedmannOld}) by $H^2$ to obtain \\begin{equation} { k c^2 \\over H^2 R^2 }  = \\Omega - 1 . \\end{equation}   Given this nomenclature, a more generalized Friedmann equation of energy can be written in a dimensionless form that explicitly incorporates all possible stable, static, isotropic energy forms described by an integer $n$. Using $H={\\dot a}/a$ in Equation~(\\ref{FriedmannOld}) and dividing each side by the square of the Hubble parameter $H_0^2$ (at $a=1$), we obtain \\begin{eqnarray} \\label{friedmanna} \\Big(\\frac{\\dot a}{H_0}\\Big)^2 = (1 - \\Omega) +\\sum_{n,m=-\\infty}^{\\infty} (\\Omega_n^+ \\, a^{2-n} + \\Omega_m^- \\, a^{2-m}) . \\end{eqnarray} Note that Eq.(\\ref{friedmanna}) can be written to highlight only present day observables by substituting $1/a = R_0/R = (1+z)/(1+z_0) = (1+z)$, where $z$ is the redshift and $z_0 = 0$ is the redshift at the epoch where $a=1$. ",
        "conclusions": "We know that our present universe is composed predominantly of positive energy and is expanding\\cite{blumenthal84,bucher99,huterer01,perlmutter98,zlatev99,spergel03}. Would it be possible for some stable negative energy form of negligible cosmological density in the present epoch but could grow to dominate and determine the future of the universe? The only way this could happen would be if some stable negative energy form with lower $n$ and $w$ than the current dominant stable energy form.  Since the dominant energy form is currently thought to be $n = 0$ ($w = -1$) dark energy, only a form of negative phantom energy could fill such a role.   Were any amount of stable negative phantom energy to exist in our universe, it would grow as the universe expands and eventually begin to dominate the universe. Since negative phantom energy is attractive such a component could stop the expansion of the universe and propel it into a contraction.  Even though the effect of this negative phantom energy would diminish and eventually become negligible, the universe would continue to contract, possibly ending in a big crunch.  Since this hypothetical stable negative phantom energy is undetectably small at present, it is really not possible to know with certainty that it does not exist, and therefore it is impossible to know the true future of our universe for this scenario.       Alternatively, would it be possible for some stable negative energy form of negligible cosmological density now but may have been significant in past of our observable universe? One  realization of this scenario would involve some stable negative energy form with higher $n$ and $w$ than the current stable positive energy form.  We know that our universe has some positive radiation and in at least two early epochs in the history of the universe, positive radiation was the cosmologically dominant energy form. Another realization demands a presence of stable negative ultralight in the current epoch. Negative ultralight energy would be gravitationally repulsive. Therefore it is possible that a previously contracting universe contained such negative ultralight, and this negative ultralight stopped the contraction and started an expansion.  Once the universe started expanding, this negative ultralight would again dissipate and again dilute to undetectable density. Therefore, since it is  not possible to confirm the existence of negative ultralight, it is not possible to know the true past of our universe for this scenario -- perhaps it started in a big bang or perhaps it underwent a big bounce.     Our universe could be an oscillating universe were it to contain even minuscule and undetectable amounts of both negative phantom energy and negative ultralight energy. As  $a$ becomes large, the negative phantom energy could suddenly begin to dominate and cause the universe to turnover and start contracting. This phantom energy would then again become effectively undetectable. As the universe contracts and its scale factor $a$ becomes small enough, the small amount of negative ultralight energy would suddenly dominate universe and, being repulsive, cause the universe to bounce. As the universe expands once again, the density of ultralight would again drop to undetectability and the cycle could repeat. Currently, it seems this type of universe cannot be ruled out theoretically\\cite{steinhardt03}."
    },
    "1402/1402.3591_arXiv.txt": {
        "abstract": "The X-ray spectra of many active galactic nuclei (AGN) exhibit a `soft excess' below 1~keV, whose physical origin remains unclear.  Diverse models have been suggested to account for it, including ionised reflection of X-rays from the inner part of the accretion disc, ionised winds/absorbers, and Comptonisation.  The ionised reflection model suggests a natural link between the prominence of the soft excess and the Compton reflection hump strength above 10~keV, but it has not been clear what hard X-ray signatures, if any, are expected from the other soft X-ray candidate models.  Additionally, it has not been possible up until recently to obtain high-quality simultaneous measurements of both soft and hard X-ray emission necessary to distinguish these models, but upcoming joint \\emph{XMM-NuSTAR} programmes provide precisely this opportunity.  In this paper, we present an extensive analysis of simulations of \\emph{XMM+NuSTAR} observations, using two candidate soft excess models as inputs, to determine whether such campaigns can disambiguate between them by using hard and soft X-ray observations in tandem.  The simulated spectra are fit with the simplest ``observer's model'' of a black body and neutral reflection to characterise the strength of the soft and hard excesses.  A plot of the strength of the hard excess against the soft excess strength provides a diagnostic plot which allows the soft excess production mechanism to be determined in individual sources and samples using current state-of-the-art and next generation hard X-ray enabled observatories. This approach can be straightforwardly extended to other candidate models for the soft excess. ",
        "introduction": "\\label{sec:intro}  Matter accreting onto supermassive black holes produces emission from radio--to--X-ray energies, with a variety of different physical processes being relevant in each band.  The direct luminous thermal power output from the central accretion disc appears in the optical--UV regime; a corona of electrons above the accretion disc Compton up-scatters this radiation to produce a power-law X-ray tail extending to hundreds of keV or greater depending on the coronal temperature.  This ensemble is often shrouded by dust and gas which introduces absorption features including a low-energy fall-off from photoelectric absorption, and troughs/edges from ionised absorption. Fall-offs are observed at higher energies due to Compton-scattering, and X-ray reflection from the accretion disc can reprocess coronal emission to produce a distinctive set of continuum features, including a reflection hump peaking at 30~keV and a broad Iron line at or near 6.4~keV.  Many AGN exhibit an excess at soft X-ray energies ($\\lesssim 2$~keV) over the coronal power-law.  This was first noticed by \\cite{1980ApJ...241L..13H} and \\cite{1981ApJ...251..501P}, and more definitively confirmed with broad-band observations by \\cite{1985ApJ...297..633S} and \\cite{1985MNRAS.217..105A}.  The feature can be modelled as a black body with an almost universal effective temperature of $\\sim$0.1--0.2~keV.  The physics of this feature remains uncertain, despite it contributing a potentially significant fraction of the total luminous power output of AGN.  The small range of effective temperatures observed disfavours a scenario where it represents the hard tail of the accretion disc spectrum peaking in the far-UV (e.g., \\citealt{2009MNRAS.394..443M,2010MNRAS.406.2591P}), as sources with different masses and accretion rates would be expected to exhibit soft excesses with different temperatures.  Instead, atomic process are often invoked to explain the soft excess, whereby a series of lines are smeared together to produce the feature.  Two mechanisms can be invoked here: 1) ionised reflection with light bending, in which the series of soft X-ray lines are relativistically blurred due to being produced very close to the centre of the accretion flow of a rapidly-spinning black hole \\citep{2005MNRAS.358..211R,2006MNRAS.365.1067C,2013MNRAS.428.2901W}, or 2) ionised absorption, where the soft excess is not really an `excess' at all, but the remnants of the true power law at low energies, with the high energy power-law being absorbed by smeared, ionised absorption with very high cloud velocities \\citep{2004MNRAS.349L...7G}.  Both models produce statistically acceptable fits to \\emph{XMM} data for the same set of PG quasars (\\citealt{2007MNRAS.381.1426M}, \\citealt{2006MNRAS.365.1067C}), but both required strong relativistic smearing of either the emission or absorption features. \\cite{2008MNRAS.386L...1S} found that this was problematic in the absorption model case due to the extreme terminal velocities required for the outflows; however, the model can still explain the data if the absorption is partially covering or clumpy.  In order to naturally produce a smooth soft excess with ionised reflection, one requires strong relativistic smearing from a rapidly spinning black hole, smoothing the soft features, and light bending from a corona with variable height above the accretion disc also allows the strength of the reflected emission (and thus of the soft excess) to vary relative to the powerlaw, and even to dominate at times. This is also supported by the observation of a time lag between the soft excess and direct continuum in the Narrow-Line Seyfert 1 galaxy 1H 0707-495 \\citep{2009Natur.459..540F,2010MNRAS.401.2419Z} and in a sample of 32 AGN \\citep{2013MNRAS.431.2441D}. Other processes can also account for the soft excess, such as Comptonisation of inner-disc photons (e.g., Done et al. 2012), and most recently, magnetic reconnection has also been suggested \\citep{2013ApJ...773...23Z}.  It is not certain, using soft-X-ray observations alone, what process is responsible for producing the soft excess (\\citealt{2006A&A...449..493C}, \\citealt{2008A&A...482..499D}, \\citealt{2013ApJ...777....2L}).   In this paper, we outline the different hard X-ray signatures that two physically-motivated models are expected to produce, and outline how using relatively simple fitting methods that do not require very high signal-to-noise ratio data, one can begin to distinguish between these models using the present and future generation of X-ray observatories.  Understanding the soft excess is important for determining the true bolometric output from accretion, and potentially provides a perspective on the extreme physical environments in AGN.  If light-bending and ionised reflection are important, it implies that the soft excess constitutes a real component of the ionising luminosity from the central engine and requires rapidly spinning black holes (e.g. \\citealt{2013MNRAS.428.2901W}); whereas if ionised absorption is more relevant in some sources, it implies that the soft `excess' belies a much larger amount of ionising flux which we do not see due to high-energy absorption.  These two scenarios can drastically change the inferred bolometric luminosity from the central engine.  The implications for the bolometric luminosity from other models are less clear, highlighting the pressing need to understand how the soft excess is produced.  In the ionised reflection model for the soft excess, we expect to see an accompanying hard ($>$10~keV) excess, but simultaneous measurements of the soft and hard bands exist only for a few sources.  Hard X-ray observatories such as \\emph{RXTE} and \\emph{Swift/BAT} have revealed such hard excesses in many AGN, typically a smooth feature that can be produced by reflection from either the accretion disc (ionised) or more distant neutral material (e.g. the inner edge of the putative `torus', or in general, clouds of absorbing gas surrounding the AGN).  The overall strength of this `Compton hump' can be measured simply using the reflection parameter $R$ of the {\\sc pexrav} model \\cite{1995MNRAS.273..837M} available in the X-ray spectral fitting package {\\sc xspec} \\citep{1996ASPC..101...17A}. The value of $R$ in this model represents the total contribution from both distant and ionised reflection \\citep{2007MNRAS.382..194N,2010MNRAS.408..601W,2013ApJ...772..114R}, although complex, Compton-thick absorbers may also be able to explain the hard excess \\citep{2013ApJ...762...80T,2013ApJ...773L...5M,2009ApJ...698...99T}.  Disentangling the components from ionised and neutral reflection (e.g. \\citealt{2007MNRAS.382..194N}) is important in trying to understand the soft excess, since it can only be produced by heavily blurred ionised reflection, not reflection from neutral material.  However, this requires very good signal-to-noise ratio data and is only possible for a few tens of sources currently.  To add confusion, ionised absorption can also produce a rising spectrum above $\\sim$7~keV until about 20~keV similar to that seen from a Compton reflection hump.  It is only with the recent advent of \\emph{NuSTAR} (launched 2012, \\citealt{2013ApJ...770..103H}), that we have real prospects to break the degeneracy between these different models, as do future hard X-ray sensitive missions such as \\emph{ASTRO-H} \\citep{2012SPIE.8443E..1ZT} and \\emph{ASTROSAT} \\citep{2013IJMPD..2241009P}.  In this paper, we ask whether the current state-of-the-art \\emph{NuSTAR} mission, with good sensitivity up to $\\sim$50~keV, in combination with the excellent 0.4--10~keV sensitivity of \\emph{XMM}, can be used to distinguish between these two scenarios using only simple characterisations of the soft and hard excess strengths, using only short ($\\sim$10~ks), inexpensive `snapshot' observations.  This paper is organised as follows: in section \\ref{bathints} we describe preliminary evidence for a link between soft excess and hard excess strengths; in section \\ref{sec:simulations} we describe the simulations of this relationship using ionised reflection and absorption models; in section \\ref{sec:discussion} we discuss the results of those simulations and in section \\ref{sec:summary} we summarise and discuss the implications of our findings. ",
        "conclusions": "\\label{sec:discussion}  At the outset, we note that the ``Observer's Model'' produces a successful fit to all of the simulated spectra using ionised reflection as the input model, but only 86\\% of the ionised absorption simulated spectra could be fit with this model.  This indicates that while ionised absorption can produce a wide range of spectral shapes that can mimic the appearance of a soft black body component with a hard X-ray reflection hump, a substantial minority do not.  Sources in those ranges of parameter space would exhibit spectral shapes clearly distinguishable from ionised reflection. We discuss this class of spectra further in \\S\\ref{subsec:failedabsfits}.  The two models compared here show some clearly distinct behaviour in $R-S_{\\rm softex}$ space; the results for both models are plotted in Fig.~\\ref{fig:refl_vs_softex_bothmodels_summary}.  Notably, in the ionised reflection scenario, stronger soft excesses can be produced and these are accompanied by stronger measured reflection fractions, particularly for $S_{\\rm softex} \\gtrsim 1$, $R \\gtrsim 1$.  Ionised reflection as a mechanism for the soft excess could be distinguishable from distant reflection by a more pronounced hard excess (values of $R \\gtrsim 2$) that one expects to obtain from such a physical process, coupled with a trend towards higher $R$ at higher $S_{\\rm softex}$, which one would expect to observe in large samples of AGN.  This would imply that in sources where both 1) strong soft excesses and 2) strong reflection are measured, ionised reflection is the most likely candidate model.  For sources with $0.5 \\lesssim R \\lesssim 3$ and $S_{\\rm softex} \\lesssim 1$, the two models are not immediately distinguishable.  Ionised absorption, on the other hand, can produce hard excesses with strengths $R$ that are not easily distinguishable from distant (torus) reflection even if a strong soft excess is present.  This would imply that for sources with 1) a strong soft excess ($0.3<S_{\\rm softex}<1$) but 2) a not particularly strong hard excess ($R \\lesssim 1$), ionised absorption may be a candidate model to explain the soft excess, but it could also be due to an altogether different physical process, with an unrelated hard excess due to distant, neutral reflection.  \\cite{2006A&A...449..493C} found that the strongest soft excesses can be produced by ionised absorption, not reflection, contrary to our findings here.  They investigated the different absorber conditions required to reproduce observed soft excesses in detail, alongside a blurred reflection model akin to the one we use here but with one key difference: the reflection spectrum strength is constrained to be sufficiently lower than that of the primary continuum such that prominent soft excesses cannot be seen.  In our model, the ratio of the direct {\\sc pexrav} continuum to the {\\sc reflionx} continuum is widely variable to account for light-bending effects (e.g., the coronal height varying above the accretion disc), and when $N_{\\rm pexrav}/N_{\\rm reflionx}\\sim 1$, strong soft excesses can be observed, as seen in reflection-dominated epochs of sources such as NGC 4051 \\citep{2006MNRAS.368..903P}.  Full consideration of the effects of light-bending at the inner parts of the accretion flow allows the reflection spectrum to dominate and produce such strong excesses.  Some very interesting behaviour is seen at the extrema of Fig.~\\ref{fig:refl_vs_softex_bothmodels_summary}.  Both ionised reflection and absorption allow for non-detected soft excesses alongside moderate reflection ($0.2 < R < 3.0$), but ionised absorption allows for a large range of soft excess strengths alongside negligible or no measured hard excess.  Therefore, in sources where the soft excess is prominent but reflection is undetectable, it may be appropriate to consider partially-covering ionised absorption or other non-reflection based models (e.g., Comptonisation). Finally, the region of the plot showing strong soft excesses $S_{\\rm softex}>1$ and relatively weak hard excesses $R<1$ is not populated by either of the two models considered here.  Further work needs to be done to explore if other models can occupy this part of parameter space.  We can apply these diagnostics now to the objects in Fig.~\\ref{reflvssoftex_BAT58monAGN}, although with the caveat that those results were obtained using non-simultaneous hard- and soft-X-ray data, without the overlap between the soft and hard band required to constrain reflection well.  Nevertheless, one can suggest (subject to further investigation with better data) that those objects in which reflection and soft excess strength seem to be increasing together likely have a strong contribution from ionised reflection.  NGC 4051 has been studied in depth and it is known that ionised reflection can be fit to the detailed 0.4--10~keV spectrum (\\citealt{2006MNRAS.368..903P}, \\citealt{2013MNRAS.435.1511A}).  Reflection has been suggested for Mrk 766 also \\citep{2011MNRAS.416L..94E}; but complex, multi-layered absorbers (perhaps a disk wind), and occultation by absorbing clouds have also been invoked \\citep{2007A&A...475..121T,2011MNRAS.410.1027R}.  Both ionised reflection and absorption can explain the spectrum of Mrk 841 \\citep{2011A&A...535A.113C}. The object Mrk 817, which has negligible reflection but a measurable soft excess, may well require a different model such as ionised absorption to account for its soft features.  Indeed, \\cite{2011ApJ...728...28W} mention an epoch in this source where absorbing winds were detected in UV spectroscopy of this source from 1997 and 2009, albeit without the X-ray edges due to Oxygen expected at 0.73 and 0.87~keV from such outflows.  One potential explanation may be that the $>2$~keV X-ray continuum in this source is absorbed by a highly relativistic ionised absorber, which may tally with the absorption signatures in the UV.  Again, data that extends into the $>10$~keV band would provide more definitive answers, stressing the utility of the approach outlined in this paper.  \\begin{figure} \\centerline{ \\includegraphics[width=9.0cm]{Rvssoftex_ionisedreflANDabs_contours.eps} } \\caption{Contours of reflection vs. soft excess strength from both ionised reflection and ionised absorption models, for comparison. \\label{fig:refl_vs_softex_bothmodels_summary}} \\end{figure}  Two Narrow-Line Seyfert 1 nuclei with strong soft excesses have recently been studied in detail using \\emph{XMM} data: 1H 0707-495 and IRAS 13224-3809.  \\cite{2012MNRAS.419..116F} and \\cite{2012MNRAS.422.1914D} find that reflection can successfully fit the spectrum of 1H 0707-495; \\cite{2013MNRAS.430.1408K} and \\cite{2013MNRAS.429.2917F} find the same for IRAS 13224-3809.  We use the archival \\emph{XMM} data to estimate $S_{\\rm softex}$ for both of these sources and find that they are 3.3 (1H 0707-495) and 2.6 (IRAS 13224-3809); therefore according to the scheme found in this paper, their soft excesses are both sufficiently strong to favour a reflection-dominated scenario.  Broad-band observations with \\emph{NuSTAR} should be able to confirm this and locate both objects on the $R-S_{\\rm softex}$ plot.  \\subsection{Extreme values of the hard excess strength, $R \\gtrsim 100$}  For both models tested in this paper, there are a cluster of points at $R\\gtrsim 100$, which have not so far been observed in the real AGN population.  For the ionised reflection case, further investigation reveals that all of the simulated spectra that produce $R>100$ have $N_{\\rm pexrav}/N_{\\rm reflionx} < 10$, and all of the very extreme $R$ values (i.e. $R > 200$) measured from `good' fits have $N_{\\rm pexrav}/N_{\\rm reflionx} = 1$, the lowest value of the ratio included in the simulations, corresponding to the most reflection-dominated case (see Fig.~\\ref{fig:ionisedreflsplitbyNratioandxi}, top panel).  This suggests that for reflection-dominated spectra, the spectral shapes produced can genuinely give rise to very high, `anomalous' $R$ values (under the fitting assumptions adopted in this study), depending on the spread of other intrinsic parameters i.e. $A_{\\rm Fe}$, $\\xi$ and $\\Gamma$.  We then focus only on the subset of objects with reflection-dominated spectra ($1 < N_{\\rm pexrav}/N_{\\rm reflionx} < 10$, Fig.~\\ref{fig:ionisedreflsplitbyNratioandxi}, lower panel).   We split the results into different bins of $A_{\\rm Fe}$, $\\xi$ and $\\Gamma$, selecting 3 bins for each (using logarithmic spacing for $\\xi$).  We find that the ionisation parameter $\\xi$ produces the greatest variation in measured $R$.  The lowest ionisation parameters, $1.0 < \\xi < 10.0$, give rise to almost all of the $R>100$ points. In conclusion, the most reflection-dominated spectra ($1.0 < N_{\\rm pexrav}/N_{\\rm reflionx} < 10.0$) coupled with the most weakly ionised reflectors ($1.0 < \\xi < 10.0$) can produce very strong hard excesses.  \\begin{figure} \\includegraphics[width=8.0cm]{ionisedrefl_splitbyreflectionstrength.eps} \\includegraphics[width=8.0cm]{ionisedrefl_highestreflstrength_splitbyxi.eps} \\caption{Measured hard excess and soft excess strength for the ionised reflection scenario, split by 1) ratio of direct-to-reflected component $N_{\\rm pexrav}/N_{\\rm reflionx}$ (top panel), and 2) ionisation parameter $\\xi$ (lower panel).  In the top panel, small and large black circles represent simulated spectra with $1<N_{\\rm pexrav}/N_{\\rm reflionx}<10$ (successful fits and 'good' fits respectively); red crosses (successful fits) and red circled crosses (good fits) represent $10<N_{\\rm pexrav}/N_{\\rm reflionx}<100$ spectra, and green X symbols (successful fits) and green X symbols within squares (good fits) represent $100<N_{\\rm pexrav}/N_{\\rm reflionx}<1000$ spectra.  In the lower panel, the same sequence of symbol combinations is used to split the reflection-dominated ($1<N_{\\rm pexrav}/N_{\\rm reflionx}<10$) subset of spectra into groups with $1<\\xi<10$, $10<\\xi<100$ and $100<\\xi<1000$, respectively.  The colours of the upper/lower limit arrows matches the colours of the points for the corresponding parameter groups, and the size of the upper/lower limit symbols increases going from small-to-high values of $N_{\\rm pexrav}/N_{\\rm reflionx}$ and $\\xi$. \\label{fig:ionisedreflsplitbyNratioandxi}} \\end{figure}  This range of $R$ values has not been seen in the real AGN population so far, but only a handful of observations currently exist of reflection-dominated states of AGN (e.g., \\citealt{2008MNRAS.391.2003Z, 2012MNRAS.419..116F}). Such prominent hard excesses may be found with new \\emph{NuSTAR} observations. Since the ionisation parameter is proportional to the luminosity, the very low luminosity sources (with potential for the most prominent hard excesses) may also be selected out of most X-ray surveys.  For the ionised absorption model, the highest values of the measured $R$ occur for $\\rm log(\\xi) < 2.75$.  From inspection of the spectra, we find that many of these low ionisation absorbers have absorption troughs that look more like extended `edges'.  The hard portion of the spectrum can then be fit as a pure reflection component in some cases, leading to high values of $R$.  In the remainder of low-ionisation absorber spectra, the slope of the intermediate (3---10~keV) region is too extreme to be fit by the {\\sc pexrav} component and the fit fails altogether (discussed further in \\S\\ref{subsec:failedabsfits}).  We also consider the possibility that the low exposure time (10ks) and resulting poor-quality data at high energies could lead to such anomalous $R$ values.  However, increasing the observation time to 20ks (for both ionised reflection and absorption models, Fig.~\\ref{fig:20ks}) did not produce any significant change in the results and the broad trends seen.  \\begin{figure} \\includegraphics[width=9.0cm]{RvsS_20kstest.eps} \\caption{Simulation results for the $R-S$ diagram with an observation time of 20~ks.  Key as for Fig.~\\ref{fig:refl_vs_softex_ionisedreflsimulations}. \\label{fig:20ks}} \\end{figure} \\subsection{Failed fits with ionised absorption} \\label{subsec:failedabsfits}  As mentioned in \\S\\ref{subsec:ionabs}, 14\\% of the simulated ionised absorption spectra result cannot be fit with the Observer's Model.  We investigate this class of objects in more detail. We find that the only weak discriminant of whether a spectrum fits successfully or not is the ionisation parameter of the absorber, $\\rm log(\\xi)$.  Failed fits only occur for spectra simulated with $2.1<\\rm log(\\xi)<2.75$.   There are 1029 possible parameter combinations/simulated spectra in this range (i.e., 3 bins in $\\rm log(\\xi)$ and 7 bins for each other parameter), out of which 328 model combinations fail completely.  On inspection, these spectra exhibit ionised absorber troughs that look more like edges with a very steep rise in flux towards higher energies.  When manually fit with the observer's model combination, we find that the fits consistently fail.  The power-law regime of the {\\sc pexrav} component attempts to fit to the rising part of the edge below 10~keV, but the photon index required is too extreme (see Fig.~\\ref{fig:ionisedlowionis_failed}).  This would be a readily identifiable subset of spectral shapes which cannot be fit with a simple black body and reflection model, and any soft `excess' seen would clearly be distinguishable from one produced by ionised reflection.  Of the remaining 701 low-$\\rm log \\thinspace \\xi$ spectra which are successfully fit by the Observer's model, only 50\\% qualify as 'good fits' according to the dual-$\\chi^{2}$ criteria.   Inspection of these `good fits' reveals that the power-law component is able to fit the rising 1--6~keV portion of the spectrum successfully.  It is therefore not straightforward to identify a very specific part of ionised-absorber parameter space that eludes fitting with the Observer's model, but we can say that the failing of the fit is restricted to low-ionisation absorbers.  The success rate of fitting such unusual spectral shapes may also be subtly dependent on the initial conditions employed in the Observer's model, the exploration of which is beyond the scope of this paper.  \\begin{figure} \\includegraphics[width=5.0cm,angle=270]{ionisedabs_successfulfit.eps} \\includegraphics[width=5.0cm,angle=270]{ionisedabs_failedfit.eps} \\caption{Example spectra with low-ionisation ionised absorbers, showing the attempted Observer's Model fit using the {\\sc pexrav+bbody} model combination. \\emph{Top panel}: example of a successful fit.  The input parameters for the simulated spectrum are $\\Gamma=2.0$, $N_{\\rm H}=5.0 \\times 10^{23} \\rm cm^{-2}$, $\\rm log \\thinspace \\xi=2.73$, $\\sigma = 0.43$.  \\emph{Lower panel}: example of a failed fit.  The input parameters are $\\Gamma=2.0$, $N_{\\rm H}=1.87 \\times 10^{23} \\rm cm^{-2}$, $\\rm log \\thinspace \\xi=2.42$, $\\sigma = 0.23$. \\label{fig:ionisedlowionis_failed}} \\end{figure}   \\subsection{Predictions for other comparable instruments}  We have presented results for the \\emph{XMM-NuSTAR} instrument combination since \\emph{NuSTAR} is newly launched and this instrument combination is already being used to obtain simultaneous broad-band X-ray data on AGN (e.g., \\citealt{2013Natur.494..449R,2014arXiv1401.5235M}); therefore it is the most relevant prediction for the current circumstances with real near-term prospects of verifying the work in this paper.  \\emph{ASTRO-H} and \\emph{ASTROSAT} are also on the horizon, and to check whether our predictions hold for other instruments, we also perform simulations for the ionised reflection model using the current \\emph{ASTRO-H} predicted response matrices, assuming the same 10~ks exposure time.   We caution that these responses may be more uncertain than the \\emph{XMM+NuSTAR} responses, as real responses are likely to differ significantly from the pre-launch predicted ones.  The resulting $R-S$ diagram does not show any significant difference to that obtained for \\emph{XMM-NuSTAR} (Fig.~\\ref{fig:astroh}), suggesting that the conclusions in this paper should hold for other similarly-equipped future broad-band X-ray observatories.   \\begin{figure} \\includegraphics[width=9.0cm]{RvsS_ionrefl_ASTROHtest.eps} \\caption{Results for ionised reflection using the \\emph{ASTRO-H} response matrices.  Key as for Fig.~\\ref{fig:refl_vs_softex_ionisedreflsimulations}. \\label{fig:astroh}} \\end{figure}   \\subsection{On the feasibility of observing these trends in real AGN samples}  The simulations here outline the trends expected in $R-S_{\\rm softex}$ space for two different models using a large grid of 2401 parameter combinations.  AGN samples are typically much smaller due to the competition for observation time, so we require a measure of how feasible it will be to detect these trends in real samples of AGN.  We perform further simulations to estimate the typical AGN sample size required, before a correlation (or the absence of one) can be seen between $R$ and $S_{\\rm softex}$.  We use a Monte-Carlo rejection method to simulate $N$ values of $R$ and $S_{\\rm softex}$ (corresponding to a sample of $N$ AGN) randomly determined throughout the available $R-S_{\\rm softex}$ space, using the contours determined from Figs.~\\ref{fig:refl_vs_softex_ionisedreflsimulations} and \\ref{fig:refl_vs_softex_ionisedabssimulations} as the probability distribution with which to distribute the points in $R-S_{\\rm softex}$ space.  We vary $N$ and measure the resultant Kendall's $\\tau$ correlation coefficient of the faked sample to determine when the strength of the correlation approaches that seen in the full simulations.  We do this for both ionised reflection and ionised absorption.  The variation of the correlation coefficient with $N$ for each input model is shown in Fig.~\\ref{fig:kendalltau_vs_samplesize}.  \\begin{figure} \\centerline{ \\includegraphics[width=9.0cm]{kendalltau_vs_samplesize.eps} } \\caption{Kendall's $\\tau$ correlation coefficient between $R$ and $S_{\\rm softex}$ against sample size, for a simulated AGN sample of size $N$.  The solid (black) and dashed (red) lines show the correlation coefficients measured for the whole set of simulations (corresponding to $N > 1500$), for ionised reflection and ionised absorption respectively.  The error bar shows the standard deviation from 100 Monte-Carlo realisations of such samples of $N$ objects, and indicates that as $N$ approaches $\\sim 60$, the correlations exhibited by samples of predominantly reflection-dominated or absorption-dominated samples can be clearly distinguished. \\label{fig:kendalltau_vs_samplesize}} \\end{figure}   We see that at $N \\gtrsim 60$, the uncertainty in the correlation coefficients produced by the two processes drops such that the two processes can be distinguished at a 2--3$\\sigma$ level.  Therefore, even though the physical process may be ambigious for an individual object based on its location in the $R-S_{\\rm softex}$ plane, the trend exhibited by a well-selected, representative AGN sample in this plane can provide an indication of the physics most relevant for the majority of AGN in the sample.   \\subsection{Outstanding issues}  We outline some issues to be further explored in future work.  Firstly, the ratio $N_{\\rm pexrav}/N_{\\rm reflionx}$ is currently used as an estimator of the degree of light bending in the reflection scenario.  A more physical understanding of this parameter is needed to place appropriate upper and lower limits on it for the simulations.  It is easy to conceive of a situation where the direct power-law dominates; however understanding the lower-limit on the range of physically viable ratios (here assumed to be Unity) is more complex and requires a fuller consideration of the energetics of the reflection-dominated state than is undertaken here.  This may impact the strongest observable strengths of the soft excess from the reflection scenario, as initially investigated by \\cite{2006A&A...449..493C}, but the observation of AGN in fully reflection-dominated states does support the possibility that $N_{\\rm pexrav}/N_{\\rm reflionx}$ can take very low values (even $<1$).  We have assumed uniform or log-uniform distributions for the input parameters in these simulations, as the simplest possible scenario.  However, it may be the case that the real underlying distributions are not uniform, or that there are correlations between the input parameters.  There are a handful of studies on this in the literature, presenting the photon index and luminosity distributions in AGN (e.g., \\citealt{2011A&A...530A..42C}, \\citealt{2013ApJ...763..111V}).  Previous works (e.g.,\\citealt{2008ApJ...675.1048R,2011ApJ...734..112B}) have found a range of ionisation parameters for ionised reflection consistent with that assumed here.  There are suggestions of a weak correlation between photon index and luminosity (e.g., \\citealt{2008AJ....135.1505S}) hence potential for a correlation between $\\Gamma$ and $\\xi$.  However, there is no detailed work on the true underlying distributions of physical parameters of reflectors.  It is possible that the observed trends highlighted in this work in the $R-S$ plane could be changed, and the density of points in different parts of the plot could be altered, if different underlying distributions are used.  To allow for this to be investigated in future, we make our simulation results public with the online data accompanying this paper.  The interested researcher can then draw from the provided results in a non-uniform way using more updated distributions or correlations between parameters, as and when such details become more well-constrained.  However, one instinctively expects some degree of correlation between hard- and soft-excess strengths in the ionised reflection scenario, regardless of the precise underlying parameter distributions.  One of the latest models to be proposed for the soft excess is the {\\sc optxagnf} model \\citep{2012MNRAS.420.1848D}.  Their model combines disk emission, Comptonisation and a power-law in an energetically self-consistent way, where the inner part of the accretion flow below a \\emph{coronal radius} is Comptonised to produce the soft excess. Comptonisation is the engine behind the soft excess in this model, and one does not expect it to exert any influence on the hard X-ray emission based on simple models (e.g., \\citealt{2004MNRAS.352..523P}) providing a purely standalone component for the soft excess.  Therefore, one does not expect any link between the soft excess and hard excess in this scenario.  However, this may not be the case for {\\sc optxagnf} where the disc and corona geometries are linked by energetic considerations, and we are currently undertaking a study of this model in detail, with a view to produce a similar $R-S_{\\rm softex}$ diagnostic plot to compare it with the models in this paper.  The {\\sc optxagnf} model is able to fully reproduce the optical--to--X-ray SED up to 10~keV, as shown by the comprehensive study of \\cite{2012MNRAS.420.1825J}, but its hard X-ray signatures have not been studied.  Additionally, distant reflection needs to be added in a consistent fashion if we want to compare it with the models studied here, where {\\sc pexrav} was used to provide the direct continuum along with the distant reflection.  How to do this is not clear, and since {\\sc optxagnf} has many more model parameters to consider than the models in this paper, we defer this study to a future paper (in prep.).    This work points to a scheme whereby different candidate models for the soft excess can be distinguished in a plot of measured reflection against soft-excess strength, assuming the simplest possible {\\sc pexrav + bbody} model combination is fit to the spectra, according to the scheme in Fig.~\\ref{fig:refl_vs_softex_bothmodels_summary}.  This methodology can readily be extended to other candidate models, and we are currently in the process of producing such a diagnostic for the {\\sc optxagnf} model, where photons from the inner part of the accretion disc are Comptonised to produce the soft excess \\citep{2012MNRAS.420.1848D}.  A key advantage of this approach is its economy: data of moderate quality can be used, gathered using short exposures (e.g. 10~ks in both \\emph{XMM} and \\emph{NuSTAR} for a source of the brightness of NGC 4051) to gain real physical insight into the energy production mechanisms in AGN, \\emph{without} requiring fitting of more complex models to long-exposure, very high signal-to-noise ratio data.  This approach will therefore be useful in constraining the soft excess mechanism in samples of AGN, where it may be challenging to obtain such long exposures on each source.  This approach is particularly useful for samples of AGN where trends can be discerned, although it can be used for individual AGN as well, if they lie in unambigious regions of $R-S_{\\rm softex}$ space.  Notably, ionised reflection predicts a clear relation between $R$ and $S_{\\rm softex}$, but ionised absorption does not.   As simultaneous or broad-band X-ray data comes in from \\emph{NuSTAR}+\\emph{XMM}, \\emph{ASTRO-H} and \\emph{ASTROSAT}, this plot can be populated with accurate, simultaneous determinations of the strengths of the hard and soft excesses in samples of real AGN, updating the work presented in Fig.~\\ref{reflvssoftex_BAT58monAGN}.  Using the contours presented in Fig.~\\ref{fig:refl_vs_softex_bothmodels_summary} as probability contours, we simulate a smaller sample of AGN using a Monte-Carlo Rejection method, and estimate that $\\sim 60$ AGN would be sufficient to verify the existence of a $R-S_{\\rm softex}$ correlation of comparable strengths to those found in our original simulations (Figs~\\ref{fig:refl_vs_softex_ionisedreflsimulations} and \\ref{fig:refl_vs_softex_ionisedabssimulations}).  This would amount to an easily achievable \\emph{XMM}-\\emph{NuSTAR} campaign of $\\sim$600~ks.  It will be easier to narrow down the mechanism responsible for producing the soft excess in any given source, if it is first fit with the simplest possible ``observer's model'' outlined here to locate it on the $R-S_{\\rm softex}$ plot.  Additionally, the regions of this plot occupied by large samples of AGN will also provide an indication of the most likely soft excess production mechanisms in the AGN population as a whole.  The general simulation methodology adopted here also has much potential for distinguishing between competing models in other areas of both AGN science and other fields of astrophysics."
    },
    "1402/1402.3558_arXiv.txt": {
        "abstract": "Long wavelength sonic oscillations are observed or inferred in many of the small, dark molecular clouds, the starless cores, that are the precursors to protostars. The oscillations provide significant internal energy and the time scale for their dissipation may control the rate of star formation in starless cores. Despite their importance, their origin is unknown. We explore one hypothesis that the oscillations develop as a  starless core forms from a filament. We model this process with a numerical hydrodynamic simulation and generate synthetic molecular line observations with a radiative transfer simulation. Comparison with actual observations shows general agreement suggesting the proposed mechanism is viable. ",
        "introduction": "In this paper we explore an hypothesis that the kinetic energy associated with the sonic oscillations  observed in starless cores comes from the gravitational energy released as a core forms from a filament by gravitational contraction.  The starless cores are dense concentrations of a few solar masses of molecular gas and are commonly observed in low-mass star-forming regions such as Taurus and Ophiucus \\citep{DiFrancesco2007,Andre2013}. They are significant as the future birthplaces of stars of  one-half to a few solar masses, the most  common stars in the galaxy \\citep{MyersBenson1983, Beichman1986}. Their internal oscillations are significant because they may control the evolution toward star formation. In star-forming regions such as Taurus, many starless cores are massive and dense enough that their internal thermal energy provides only about 70\\% of the energy needed for stability against gravitational collapse \\citep{DickmanClemens1983,MyersBenson1983,Evans2001,Kandori2005, Kirk2005}. The remaining energy required for stability may be in the kinetic energy of the oscillations \\citep{Keto2006,Broderick2010}.  If oscillations support the cores, then the dissipation of the oscillations by mode-mode coupling and radiation should result in gravitational collapse and star formation if the oscillation energy is not resupplied \\citep{Broderick2010}. In this case, the core lifetimes would be equal to the dissipation time \\citep{Broderick2008}. This deduction is supported by some observations that suggest that more evolved cores, those with newly formed stars, have less turbulent energy than the starless cores \\citep{Tachihara2002}.  Statistical arguments based on observation suggest that most cores have lifetimes longer than their free-fall times \\citep{Beichman1986,Jessop2000,Ward-Thompson2007}. If many cores require additional energy for stability, and the oscillations are the source of that energy, then the core lifetimes suggest that many cores are oscillating, not just the few observed with complex modes \\citep{Lada2003,Redman2006,Aguti2007}.  Despite the possible significance of the oscillations in the evolution of star-forming cores, the origin of the oscillations and the source of their energy is unknown. \\citet{KetoField2005} supposed that the oscillations were the expression of the subsonic turbulence at the bottom of the supersonic turbulent cascade that characterizes the velocities in the larger-scale interstellar medium \\citep{FBK2008}. However, they did not specify further. Similarly, \\citet{Broderick2010} assumed an initial turbulent spectrum, either Kolmogoroff or flat, but again without specfiying the origin.  The origin of the oscillations may be in the process of core formation from filaments.  As a filament fragments and contracts to form a spherical core, the gravitational energy released is converted into kinetic energy in oscillations. This hypothesis is motivated by both observation and theory. Observations of low-mass star-forming regions show that the cores are often associated with filaments \\citep{Goldsmith2008,Myers2009}. Theoretically, we know that self-gravitating filaments are unstable to fragmentation and the formation of cores. It may be that most cores form out of filaments.  We model the process of core formation from the contraction of a filamentary fragment with a hydrodynamic simulation. We follow the evolution and at various time we simulate observable spectral lines that can be compared with actual observations. We find that in the initial contraction, the velocities are mostly linear, along the filament, in the direction of the contraction. As the gas rebounds off its own thermal pressure, the fragment flattens and expands perpendicular to the length of the initial filament.  This is followed by a long period of oscillations of slowly decreasing amplitude. Eventually the velocities become more and more radial as the fragment becomes more and more spherical.   The profiles of the simulated spectral lines are generally consistent with those observed. Our conclusion is that this is a plausible hypothesis for the generation of the oscillations in starless cores. ",
        "conclusions": "A hydrodynamic simulation of a contracting filamentary fragment is able to produce oscillating cores with some complexity despite the high degree of symmetry in our choice of initial state. The general agreement between simulated and actual spectral line observations maintains the viability of this mechanism for the origin of the oscillations.    \\clearpage  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fnewxy-2p157} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fnewyz-2p157} \\end{array} $ \\caption{ Densities and velocities on slices through two orientations at 2.2 Myr, ({\\it left}) in the XY plane at Z=0, and ({\\it right}) in the YZ planet X=0. The filament was originally aligned with the X-axis. The axes are labeled in units of $10^{17}$ cm and the density in units of $2\\times 10^{-18}$ g cm$^{-3}$. The longest arrow, the largest velocity is 0.62 km s$^{-1}$. } \\label{fig:figxy-2p159} \\end{figure*}   \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig2Lnew} \\includegraphics[trim=0.10in 1.6in 0.8in 0.4in, clip,width=3.25in,angle=180]{fig2R} \\end{array} $ \\caption{ ({\\it Left}) First moment of the simulated HCO$^+$(1-0) line on a projection along the Z axis. The axes are labeled in arc minutes assuming a distance of 250 pc. Velocities in km s$^{-1}$. ({\\it Right}) Spectra at intervals of  0.42 arc minutes. The velocity axis spans a range of -1 to +1 km s$^{-1}$. The maximum brightness is 3.5 K. } \\label{fig:2p159_00_90_hcoplus} \\end{figure*}  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig3Lnew} \\includegraphics[trim=0.10in 1.6in 0.8in 0.4in, clip,width=3.25in,angle=180]{fig3R} \\end{array} $ \\caption{ Same as figure \\ref{fig:2p159_00_90_hcoplus} but on a projection on the YZ axis. The maximum brightness is 4.6 K. } \\label{fig:2p159_90_00_hcoplus} \\end{figure*}  \\clearpage  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fig4Lnew} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fig4Rnew} \\end{array} $ \\caption{ Same as figure \\ref{fig:figxy-2p159} except for a time of 2.9 Myr. The maximum velocity is 0.17 km s$^{-1}$. } \\label{fig:figxy-2p867} \\end{figure*}  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig5Lnew} \\includegraphics[trim=0.10in 1.6in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig5R}  \\end{array} $ \\caption{ Same as figure \\ref{fig:2p159_00_90_hcoplus} but for a time of 2.7 Myr. The maximum brightness in the spectra is 0.5 K. } \\label{fig:2p867_00_90_hcoplus} \\end{figure*}  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig6Lnew} \\includegraphics[trim=0.10in 1.7in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig6R} \\end{array} $ \\caption{ Same as figure \\ref{fig:2p867_00_90_hcoplus} but on a projection on the YZ axis. The maximum brightness is 0.6 K } \\label{fig:2p867_90_00_hcoplus} \\end{figure*}  \\clearpage  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fig7Lnew} \\includegraphics[trim=0.10in 5.0in 0.8in 0.5in, clip,width=3.25in]{fig7Rnew} \\end{array} $ \\caption{ (figxy-3p448) Same as figure \\ref{fig:figxy-2p159} except for a time of 3.4 Myr. The maximum velocity is 0.08 km s$^{-1}$. } \\label{fig:figxy-3p448} \\end{figure*}  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig8Lnew} \\includegraphics[trim=0.10in 1.7in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig8R} \\end{array} $ \\caption{ Same as figure \\ref{fig:2p159_00_90_hcoplus} but for a time of 3.5 Myr. The maximum brightness in the spectra is 0.8 K. } \\label{fig:3p448_00_90_hcoplus} \\end{figure*}  \\begin{figure*} $ \\begin{array}{cc} \\includegraphics[trim=0.10in 5.5in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig9Lnew} \\includegraphics[trim=0.10in 1.7in 0.8in 1.0in, clip,width=3.25in,angle=180]{fig9R}  \\end{array} $ \\caption{ Same as figure \\ref{fig:3p448_00_90_hcoplus} but on a projection on the YZ axis. The maximum brightness is 0.8 K. } \\label{fig:3p448_90_00_hcoplus} \\end{figure*}  \\clearpage  \\begin{figure} \\includegraphics[trim=3.7in 3.8in 0.2in 0.2in, clip,width=3.25in,angle=0,scale=0.8]{evolution} \\caption{ Average velocity ({\\it black solid line})and maximum density ({\\it red dotted line}) as a function of time. The average velocity is defined by the total kinetic energy, $E_{kin}$  divided by the total mass, $M$, $v=\\sqrt{2*E_{kin}/M}$. The units of density are $2\\times 10^{-18}$ g cm$^{-3}$. The figure shows the 2 periods of oscillation, longitudinal and radial with the highest densities occurring at the point of maximum compression when the velocities are low and reversing. } \\label{fig:evolution} \\end{figure}       \\noindent{\\bf Acknowledgements} A.B. thanks the Harvard-Smithsonian  Center for Astrophysics for their hospitality during multiple visits. The research of A.B. is supported by the priority program 1573 \"Physics of the Interstellar Medium\" of the German Science Foundation."
    },
    "1402/1402.6182_arXiv.txt": {
        "abstract": "A search for cosmic neutrino sources using six years of data collected by the ANTARES neutrino telescope has been performed. Clusters of muon neutrinos over the expected atmospheric background have been looked for. No clear signal has been found. The most signal-like accumulation of events is located at equatorial coordinates RA=$-$46.8$^{\\circ}$ and Dec=$-$64.9$^{\\circ}$ and corresponds to a 2.2$\\sigma$ background fluctuation. In addition, upper limits on the flux normalization of an E$^{-2}$ muon neutrino energy spectrum have been set for 50 pre-selected astrophysical objects. Finally, motivated by an accumulation of 7 events relatively close to the Galactic Centre in the recently reported neutrino sample of the IceCube telescope, a search for point sources in a broad region around this accumulation has been carried out. No indication of a neutrino signal has been found in the ANTARES data and upper limits on the flux normalization of an E$^{-2}$ energy spectrum of neutrinos from point sources in that region have been set. The 90\\% confidence level upper limits on the muon neutrino flux normalization vary between 3.5 and 5.1$\\times$\\textrm{10$^{-8}$ GeV$\\,$cm$^{-2}$s$^{-1}$}, depending on the exact location of the source. ",
        "introduction": "The scientific motivation of neutrino telescopes relies on the unique properties of neutrinos, which can be used to observe and study the high-energy Universe. Cosmic rays or high-energy photons have intrinsic limitations: the mean free path of gamma-ray photons strongly depends on their energy, while magnetic fields deflect cosmic rays, diluting the information about their origin. Neutrinos are stable, neutral and weakly interacting particles, and therefore they point directly back to their origin. In addition, neutrinos are expected to originate at the same locations where the acceleration of cosmic rays and the associated production of high-energy photons take place \\cite[]{RepProg, Stecker, Bednarek}. The first evidence of such a cosmic neutrino signal has recently been reported by IceCube ~\\cite[]{IceCube1,IceCube2}, including in particular a cluster of events close to the Galactic Centre. The better view of the Southern Hemisphere afforded by the ANTARES neutrino telescope, due to its location in the Mediterranean Sea,  provides an increased sensitivity to galactic sources of neutrinos with energies $<$ 100 TeV. This is particularly important in order to interpret the cluster of events observed by IceCube close to the Galactic Centre.  In this paper the results of the search for point sources with the data gathered between 2007 and 2012 with the ANTARES neutrino telescope are presented. After a brief description of the apparatus, the data selection and the corresponding detector performance are presented in Sections~\\ref{ANTARESdet} and~\\ref{performance}, respectively. In Section~\\ref{method}, the search method is explained. The results of the full-sky and candidate sources searches are presented in Section~\\ref{results}. The implications on some recent interpretations of the IceCube results are discussed in Section~\\ref{Implications}. Finally, the conclusions are given in Section~\\ref{conclusions}. ",
        "conclusions": "In this paper the results of a search for neutrino point sources with six years of ANTARES data (2007-2012) are presented using two complementary analyses. Firstly, a scan for point sources of the ANTARES visible sky. Secondly, a search for correlations of events with a pre-selected  list of candidate sources for neutrino emission. In the first case, the most significant cluster has a post-trial p-value of 2.7\\% (a significance of 2.2$\\sigma$). In the case of the candidate list study, the largest excess corresponds to HESS J0632+057 with a post-trial p-value of 6.1\\% (1.9$\\sigma$). Both results are compatible with a pure background hypothesis. The derived flux upper limits are the most restrictive in a significant part of the Southern sky. The possibility that the accumulation of 7 events reported by IceCube near the Galactic Centre is produced by a single point source has been excluded. These results show the potential of neutrino telescopes in the Northern hemisphere, such as the planned KM3NeT observatory \\citep[]{KM3Net}, to interpret the increasing evidence of cosmic neutrino fluxes."
    },
    "1402/1402.1594_arXiv.txt": {
        "abstract": "We present statistical evidence suggesting more than one population in the energy distributing of  pulsar glitches, which implies the presence of different mechanisms  accessing different energy ranges responsible for glitches. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0889_arXiv.txt": {
        "abstract": "Determining the presence of widely separated substellar-mass companion is crucial to understand the dynamics of inner planets in extrasolar planetary systems (e.g. to explain their high mean eccentricity as inner planets are perturbed by the Kozai mechanism). We report the results of our \\textit{Spitzer}/Infrared Array Camera (IRAC) imaging search for widely separated (10 -- 25$^{\\prime\\prime}$) substellar-mass companions for 14 planet-host stars within 15~pc of the Sun. Using deep 3.6 and 4.5~$\\mu$m observations in subarray mode, we found one object in the field of 47~UMa with [3.6]$-$[4.5] color similar to a T5 dwarf, which is, however, unlikely to share common proper motion with 47~UMa. We also found three objects with brown-dwarf-like [3.6]$-$[4.5] color limits in the fields of GJ~86, HD~160691, and GJ~581, as well as another in the field of HD~69830 for which we have excluded common proper motion. We provide model-based upper mass limits for unseen objects around all stars in our sample, with typical sensitivity to 10~M$_{J}$ objects from a projected separation of 50 to 300 au from the parent star. We also discuss our data analysis methods for point-spread-function subtraction, image co-alignment, and artifact subtraction of IRAC subarray images. ",
        "introduction": "Large bodies in the periphery of planetary systems have the potential to wreak havoc on the orbits of inner planets through secular interactions. \\citet{naoz11} posit this effect to be responsible for the migration of hot Jupiters; \\citet{marzari02} describe how planet-planet scattering can result in moderately eccentric planetary orbits; and \\citet{takeda05} demonstrate how Kozai-type interactions \\citep{kozai62} can result in highly eccentric planet orbits, inward of a large planet or brown dwarf.  Given that conventional planet formation models, including both core accretion \\citep{pollack96} and gravitational instability \\citep{boss97}, describe planets forming in mostly circular, well-aligned circumstellar disks, it is important to explain why the mean eccentricity of discovered exoplanets is currently 0.19\\footnote{http://exoplanet.eu, as of Oct 7, 2013}, with many of the highest eccentricities found in single-planet systems (78\\% with $e>0.5$ are single planets). If secular effects involving a massive companion are responsible, these companions may be observable, if they have not yet been ejected from the system. \\citet{ford08} conclude that most planet scattering instabilities should occur on timescales comparable to the planet formation timespan, but do acknowledge that triple planet systems in certain configurations could potentially be quasi-stable on timespans from $10^{6}$ to $10^{10}$ years, so it is not impossible to have planet ejection occur in an old system.  \\citet{gizis01} found brown dwarf companions $>1000$~au from the primary, and \\citet{luhman07} found HD~3651B, a brown dwarf orbiting 476~au from a star with an eccentric planet. However, the same formation models mentioned above do not explain how such a massive companion could form so far from the parent star: that far-separated substellar companions and highly eccentric planets both exist deserves explanation. One such explanation is that all objects form close to their parent stars and end up in their current configurations through scattering events (\\citealt{boss06}, \\citealt{nagasawa11}).  Whether causing havoc or not, widely separated companion brown dwarfs are interesting in their own right, as relatively few are known: of the 992\\footnotemark[1] substellar-mass companion objects discovered, only 36 (all with mass $>1~M_{J}$) are farther than 10~au from their parent, and 21 of those lie between 10-300~au. In part, this is a selection effect: the timescales necessary to find these objects with either radial velocity (RV) or transit methods are prohibitive, and all of the aforementioned brown dwarfs have been detected by direct imaging. Ground-based direct imaging searches have been successful in finding gas giant companions (e.g. \\citealt{carson13}, \\citealt{marois10}, \\citealt{chauvin05} and \\citealt{neuhauser05}), but have lower sensitivity to the coolest T- and Y-dwarfs (which have distinctive [3.6]$-$[4.5] colors) because of telluric absorption in the thermal infrared. Space telescopes get around this problem, and WISE has been successful in finding many field T- and Y-dwarfs \\citep{kirkpatrick11}, but is less sensitive to brown dwarfs around stars because of the high contrast ratios involved. Studies like \\citealt{bergfors13}, \\citealt{ginski12}, and \\citealt{roell12} have searched for stellar-mass companions to planet-host stars, but are generally not sensitive to brown-dwarf-mass objects.  The inner working angle of any direct-imaging companion search is limited by the brightness of the parent star compared to its companion. At optical wavelengths, this is determined by the contrast ratio between the parent star's emission and the companion's reflected light. This can be problematic at close angular distances or if the starlight saturates the detector. However, the contrast ratio in the infrared (e.g. at 3.6~$\\mu$m), where thermal emission dominates over reflected light, is much more manageable. This lower ratio opens up the possibility of searching for planetary mass companions with advanced Point Spread Function (PSF) subtraction techniques, provided that the PSF is both stable and well spatially sampled. For these reasons, we have used the InfraRed Array Camera (IRAC, \\citealt{fazio04}) onboard the Spitzer Space Telescope \\citep{werner04} to search for companions around 14 nearby stars. Using IRAC gives the added advantage of being able to characterize T-dwarfs, in particular, using their characteristically red [3.6]-[4.5] color due to methane absorption bands \\citep{burrows03}.  This work builds on our previous IRAC low-mass companion searches (e.g. \\citealt{marengo06}, \\citeyear{marengo09}, \\citealt{carson11}), which have resulted in the discovery of two brown dwarf companions \\citep{luhman07}, including a T7.5 companion of the exoplanet host star HD~3651. In \\citet{marengo09}, in particular, we have developed the same technique used in the work presented here, adopting IRAC shorter frame-time subarray observing mode. This allows us to reduce the primary star saturation and the area with high PSF-subtraction residuals, narrowing our inner working angle from 20$^{\\prime\\prime}$ to 5$^{\\prime\\prime}$. Building on this previous work, we used IRAC's subarray mode to probe for $>5~M_{J}$ companions at separations of 25 to 350~au (for a star 15~pc distant) around a sample of 14 nearby planet host stars.  Herein, we present our target list in Section~\\ref{targetselection}, our analysis methods (Section~\\ref{dataanalysis}), sensitivity limits (Section~\\ref{sensitivitysection}), and companion candidates selected on the basis of their IRAC colors (Section~\\ref{detectedsources}). In Section~\\ref{discussion} we discuss our results. ",
        "conclusions": "\\label{discussion}  Using the combined photometry from our subarray images and the available archival full-frame images, we found four potentially interesting candidates: one with matching brown dwarf colors and three objects with compatible [3.6]$-$[4.5] color lower limits. In absence of common proper motion confirmation, we are unable to determine if these candidates are true companions, or unrelated background quasars or mass-losing giants (see \\citealt{stern07}; \\citealt{reiter14} for typical colors of background red sources). This ambiguity could, in principle, be resolved by acquiring JHK photometry (see \\citealt{marengosanchez09}), albeit new near-IR data capable of detecting these objects would also likely provide accurate astrometry, sufficient to test for common proper motion.  We presented details of our PSF- and artifact-subtraction procedure, in which we were able to automatically co-align images within 0.0027$^{\\prime\\prime}$ for purposes of PSF subtraction and effective removal of electronic artifacts. As anticipated, subarray observations allowed a smaller inner working angle than full-frame observations of the same field of view, which were dominated by PSF subtraction residuals.  For all stars, we calculated upper limits on unseen T- or Y-dwarfs within the mass and projected semi-major axis ranges given in Figures~\\ref{sensitivity} and \\ref{sensitivity2}. Outside of 10$^{\\prime\\prime}$ from most sources, our sensitivity is in the range of 10~M$_{J}$. Our best sensitivity is for close, young stars like GJ~849 (5~M$_{J}$) and SCRJ~1849-6357 ($<5~M_{J}$). Our worst sensitivity (25~M$_{J}$) was for $\\upsilon$~And due to its brightness: artifacts overwhelmed the image.  We could not confirm the existence of any new wide planetary or brown dwarf companions around the stars in our sample. We have ruled out a section of parameter space, but there is a degeneracy for perturbers between mass and semi-major axis. It is still possible for less-massive companions closer than our inner working angle to exist (ongoing systematic searches with new ground-based adaptive optics systems will explore the remaining parameter space), or there could be some other reason for the measured high eccentricities of planets in those systems.  Follow up observations are necessary for HD~160691-8, GJ~86-1, and GJ~581-3 to obtain 3.6~$\\mu$m photometry and check for common proper motion. If the candidates are confirmed, numerical simulations will be necessary to determine whether they can account for the high eccentricities of planets in their respective systems."
    },
    "1402/1402.5654_arXiv.txt": {
        "abstract": "We present a large-scale clustering analysis of radio galaxies in the Very Large Array (VLA) Faint Images of the Radio Sky at Twenty-cm (FIRST) survey over the Galaxy And Mass Assembly (GAMA) survey area, limited to $S_{1.4 \\textrm{ GHz}}$ \\textgreater 1~mJy with spectroscopic and photometric redshift limits up to $r < 19.8$ and $r <22$~mag, respectively. For the GAMA spectroscopic matches, we present the redshift-space and projected correlation functions, the latter of which yielding a correlation length $r_0 \\sim 8.2$ $h^{-1}$Mpc and linear bias of $\\sim1.9$ at $z\\sim0.34$. Furthermore, we use the angular two-point correlation function $w(\\theta)$ to determine spatial clustering properties at higher redshifts. We find $r_0$ to increase from $\\sim6$ to $\\sim 14$ $h^{-1}$Mpc between $z=0.3$ and $z=1.55$, with the corresponding bias increasing from $\\sim 2$ to $\\sim 10$ over the same range. Our results are consistent with the bias prescription implemented in the SKADS simulations at low redshift, but exceed these predictions at $z>1$. This is indicative of an increasing (rather than fixed) halo mass and/or AGN fraction at higher redshifts or a larger typical halo mass for the more abundant FR\\textsc{I} sources. ",
        "introduction": "In recent decades, measurements of the cosmic microwave background radiation (CMB; e.g. \\citealt{komatsu11}) have proven the very early Universe to be remarkably isotropic with tiny under- and overdensities that have grown to form the vast and intricate structures we see today. Galaxies and clusters observed today are far removed from their almost homogeneous beginnings and we require large numbers of them to piece together a statistical picture on cosmological scales. Clustering measures on large scales can be used to investigate not only the relationships between galaxy populations found by various techniques probing different epochs and masses, but also broader cosmological phenomena such as baryon acoustic oscillations (BAO; e.g. \\citealt{eisenstein05, percival10, blake11}), cosmic magnification (e.g. \\citealt{scranton05, wang11}), or the integrated Sachs-Wolfe (ISW) effect (e.g. \\citealt{mcewen07, giannantonio08, raccanelli08}).  Cosmological applications require information about the gravitating mass distribution in the Universe, which in a $\\Lambda$CDM cosmology is strongly tied to the dark matter distribution. Direct observations tell us only about the baryonic matter, from which we must infer the dark matter distribution. Various tools exist for measuring the clustering signal of an observed source catalogue, such as nearest neighbour measures (e.g. \\citealt{bahcall83}), counts-in-cells (e.g. \\citealt{magliocchetti99,blake02b,yang11}), correlation functions (e.g. \\citealt{groth77,bahcall83,blake02a,blake02b,croom05}) and power spectra (e.g. \\citealt{cole05,percival07,komatsu11}). Due to its relative simplicity to calculate, and relation to its Fourier transform (the power spectrum), the two-point spatial correlation function has become a standard in quantifying cosmological structure. A means by which we can quantify the extent to which the observable and dark matter are tied using the correlation function is through the bias parameter $b(z)$. This relates the spatial correlation function (see Section \\ref{spatial}) of an observed galaxy population to that of the underlying dark matter. The bias quantifies the difference in the clustering of the dark matter haloes acting solely under gravity and of galaxies inhabiting those haloes with other effects making their structure more or less diffuse. This has a heavy dependence on the galaxy masses and the epoch under consideration (e.g. \\citealt{seljak04}).  Local galaxies with masses comparable to the Milky Way are a relatively unbiased tracer of mass, with a present day correlation length (the clustering scale at which the correlation function falls below unity; see Section \\ref{spatial}) of $r_0 = 5.4$ $h^{-1}$Mpc found in the early CfA redshift survey \\citep{davis83}. While the power law slope is consistently found to be $\\sim 2$ in most studies, large variations in $r_0$ are found depending on the populations and respective epochs being observed, from $\\sim 5$ $h^{-1}$Mpc for local galaxies to $\\sim 25$ $h^{-1}$Mpc for Abell clusters (e.g. \\citealt{bahcall83}). At the lower end of the clustering scale, \\citet{saunders92} find $r_0 \\sim 4$ $h^{-1}$Mpc for IRAS starburst galaxies and using the 2dF QSO Redshift Survey \\citet{croom01} show a scale roughly constant with redshift at $\\sim 5$ $h^{-1}$Mpc for quasars at $0 \\lesssim z \\lesssim 2.5$ (see also \\citealt{croom05,ross09,ivashchenko10}). Likewise, \\citet{kovac07}  find that $z \\sim 4.5$ Ly-$\\alpha$ emitters exhibit a relatively short clustering length of $4.6$ $h^{-1}$Mpc, consistent with their being progenitors of Milky Way type local galaxies (see \\citealt{nilsson09} and references therein), and note the similarity with Lyman break galaxies (LBGs) at $z \\sim 3.8$ and $z \\sim 4.9$, suggesting that the two populations reside in the same host halos but with a relatively low duty cycle. Local $L \\gtrsim L_{\\ast}$ ellipticals, however, are a highly clustered population with various authors finding $r_0 \\sim 7$--$12$ $h^{-1}$Mpc (e.g. \\citealt{guzzo97,willmer98,norberg02}). Similarly high clustering lengths are found for extremely red objects (EROs) at $z \\sim 1$ (e.g. \\citealt{daddi01,mccarthy01,roche02}) and distant red galaxies (DRGs) at $z \\sim$1--2 \\citep{grazian06,foucaud07}, both posited as progenitors of the local bright ellipticals. Indeed, similar results for radio galaxies (e.g. \\citealt{cress96,overzier03}) give weight to the suggestion by \\citet{willott01} that EROs and radio galaxies are identical and seen at different evolutionary stages, based on their findings of ERO-like hosts of radio galaxies in the 7C Redshift Survey.   Radio surveys are ideal for the purpose of carrying out such large-scale statistical measurements for a number of reasons. From a logistical point of view, radio wavelengths are useful because they occupy a uniquely broad atmospheric window in the electromagnetic spectrum, not being significantly absorbed by Earth's atmosphere. This allows observations to be carried out from terrestrial telescopes, reducing costs compared with launching and operating space telescopes in orbit. In a wide radio survey with relatively shallow flux density limit ($S_{1.4} \\gtrsim 1$ mJy) extragalactic sources dominate in the form of synchrotron radiation in active galaxies with supermassive black holes at their centre. By observing such powerful sources, and with the long radio wavelengths being immune to absorption by dust, the ISM and our own atmosphere, ground based telescopes are able to observe the whole sky relatively unobscured and to very high redshifts.  Of particular significance for the science discussed here is that these radio-loud active galactic nuclei (AGN) predominantly reside in massive elliptical galaxies (e.g. \\citealt{jarvis01,willott03,mclure04,herbert11}) and, as such, provide good tracers of large-scale clustering and the underlying dark matter distribution. Radio surveys have been exploited with this aim for some time, such as the 4.85 GHz 87GB \\citep{kooiman95} and Parkes-MIT-NRAO \\citep{loan97} surveys, the 325 MHz Westerbork Northern Sky Survey (WENSS; \\citealt{rengelink98}) and 1.4 GHz VLA surveys, Faint Images of the Radio Sky at Twenty-cm (FIRST; \\citealt{becker95}) and the NRAO VLA Sky Survey (NVSS; \\citealt{condon98}) (e.g. \\citealt{cress96,blake02a,overzier03, best05, fine11}), but their potential uses are limited without additional data at other wavelengths. While the large redshift range of radio surveys can be advantageous over more targeted optical surveys, this can also wash out the clustering signal due to the superposition of the influence of differently behaving subpopulations (e.g \\citealt{wilman08}). This effect can be alleviated by the use of complementary optical surveys to aid the separation of radio sources by simple redshift binning or other methods.  As technology advances, observations at radio wavelengths will be able to constrain cosmological parameters with more precision than ever \\citep{raccanelli12, camera12}, as well as provide the data for much more galactic and extragalactic science \\citep{norris11}. Moving forward, the large numbers of sources observed in various overlapping surveys will make multi-wavelength studies not only more common, but essential to reliably identify objects by their emission across the entire spectrum and delineate the clustering properties of various subpopulations of any given individual survey.  It is with this goal in mind that in this paper we combine the available radio data from the FIRST survey, and optical/near-infrared photometry and spectra from Galaxy and Mass Assembly (GAMA; \\citealt{driver11}), Sloan Digital Sky Survey (SDSS; \\citealt{york00}) and UKIRT Infrared Deep Sky Survey (UKIDSS; \\citealt{lawrence07}) to measure the large scale clustering properties of the $S_{1.4} > 1$ mJy radio emitting galaxies. We choose to use FIRST over NVSS for its marginally greater depth and considerably better resolution which is essential for meaningful cross-identification with the optical surveys. Furthermore, we employ a parent redshift distribution from the simulated radio catalogue of the SKA Design Study (SKADS; \\citealt{wilman08}) which has the express intention of acting as a test bed to inform the use of the forthcoming Square Kilometre Array.  The outline of this paper is as follows: Section \\ref{data} describes the surveys and catalogues used, while Sections \\ref{angular} and \\ref{Limber} detail the angular correlation function measurements and their deprojection to infer spatial clustering properties. Section \\ref{spatial} describes direct measurements of the spatial clustering properties through the spatial correlation function and projected correlation function and Section \\ref{massbias} describes the linear bias found from each of these methods out to $z\\sim1.5$. In Section \\ref{conclusions} we summarise our results and conclusions.  The cosmological model used throughout this paper is the flat, $\\Lambda$CDM concordance cosmology where $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$ and $\\sigma_8 = 0.8$. All distances are kept in units of $h^{-1}$Mpc where $H_0 = 100 h$ km s$^{-1}$Mpc$^{-1}$ and $h$ is not explicitly assumed. ",
        "conclusions": "Within our matched sample and incorporating simulated data, we observe the increasing of $b(z)$ in line with models corresponding to fixed masses \\citep{mo96, matarrese97}, although we find a relatively  high value for the high-redshift sources, particularly when using the linear clustering parameter. Whilst the GAMA data points do not place especially tight constraints on this evolution, and the well-constrained point from $\\Xi(\\sigma)$ is limited to one epoch, the assumption of the redshift distributions for the full and unmatched subsamples probe redshifts up to $z = 1.55$ and give more shape to the evolution. The agreement is good at lower redshifts, the GAMA-matched results (\\textit{red}) lying $<1\\sigma$ from the model prediction as well as agreeing with the 2SLAQ radio-detected LRGs examined by \\citet{wake08} ($b\\sim3.0$). The higher redshift points, however, significantly exceed the values prescribed in the SKADS simulation. The wider FIRST sample is an exception, falling significantly below the trend displayed by the narrower samples, but closely matching the Mo \\& White bias prescription from SKADS. On face value, this apparent excess clustering suggests a greater proportion of AGN observed than was assumed in the simulations, or perhaps a considerable underestimate of the typical halo masses of any or all of the galaxy types considered.  Our unmatched FIRST sources, especially, in the GAMA field are more strongly clustered than would be expected by simply subtracting the cross-matched $N(z)$ from the SKADS $N(z)$ and assuming the same population as a matching SKADS sample. This means either that the assumed redshift distribution is skewed towards low redshift, or the fraction of more massive galaxies is higher than expected. Inspection of Figure~\\ref{zdists} shows that our assumed $N(z)$ for this point contains a significant fraction of $z<0.5$ sources where one would expect to observe a greater proportion of less massive, star-forming galaxies rather than the more massive AGN. Given that if these low-redshift objects were detected in FIRST, we would expect them to have optical counterparts in SDSS/GAMA (cf. almost complete cross-identification at $z\\sim0.6$ in Figure \\ref{zdists}), it is perhaps unrealistic to assume these missing low-$z$ sources remain in our unmatched sample. The cause of this large fraction of unmatched sources is discussed in Section~\\ref{redshifts}, however to determine the effect that these sources have on our measured clustering length and bias we implement cuts on the SKADS redshift distribution, where we remove all sources at $z< 0.05$ and $z< 0.1$ from the simulations. This is a coarse way of simulating the combined effect of the resolution bias and the removal of very low-redshift sources with photometric redshifts with $z<0.002$. We then recalculate the correlation length and bias using these new redshift distributions and find that the correlation length and bias increases significantly ($\\sim$25--50 per cent increase in bias). % Although subject to many uncertainties, this may point to a higher proportion of highly biased objects in the radio source population at high redshift.  One of the main uncertainties in the SKADS simulation is indeed the evolution and bias of the FR\\textsc{I}-type objects which dominate the source counts at the flux density limit of the FIRST survey at $z>0.5$. In the SKADS simulation the FR\\textsc{I}s are less biased than the FR\\textsc{II}s, however it is becoming clear that there is a large overlap in the stellar mass distributions between the generally less radio luminous FR\\textsc{I}s and their FR\\textsc{II} counterparts at high redshifts \\citep[e.g.][]{mclure04,herbert11}, thus it is possible that the bias prescription for the FR\\textsc{I} sources in the SKADS simulation is underestimated. If we assign a similar halo mass to the FR\\textsc{I} as for the FR\\textsc{II}s then we find that the expected average bias of radio sources in our sample ($z\\sim1.21$) to be in the region of $b=5.7$, this is much more closely aligned with our measured value of $b=5.95^{+0.48}_{-0.53}$.  Deeper surveys over smaller areas will be able to address this issue better than the relatively shallow FIRST data. Indeed, deep multi-wavelength surveys are now beginning to tackle the question of redshift evolution \\citep[e.g.][]{smolcic09,mcalpine11,simpson12,mcalpine13} and with slightly more area could measure the clustering length."
    },
    "1402/1402.5238_arXiv.txt": {
        "abstract": "We present the first constraint on Faraday rotation measure (RM) at submillimeter wavelengths for the nucleus of M 87. By fitting the polarization position angles ($\\chi$) observed with the SMA at four independent frequencies around $\\sim$230 GHz and interpreting the change in $\\chi$ as a result of \\emph{external} Faraday rotation associated with accretion flow, we determine the rotation measure of the M 87 core to be between $-$7.5$\\times$10$^{5}$ and 3.4$\\times$10$^{5}$ rad/m$^{2}$. Assuming a density profile of the accretion flow that follows a power-law distribution and a magnetic field that is ordered, radial, and has equipartition strength, the limit on the rotation measure constrains the mass accretion rate $\\dot{M}$ to be below 9.2$\\times$10$^{-4}$ M$_{\\odot}$~yr$^{-1}$ at a distance of 21 Schwarzchild radii from the central black hole. This value is at least two orders of magnitude smaller than the Bondi accretion rate, suggesting significant suppression of the accretion rate in the inner region of the accretion flow. Consequently, our result disfavors the classical \\emph{advection dominated accretion flow} (ADAF) and prefers the \\emph{adiabatic inflow-outflow solution} (ADIOS) or \\emph{convection-dominated accretion flow} (CDAF) for the hot accretion flow in M 87. ",
        "introduction": "Revealing the mass accretion process onto supermassive black holes (SMBHs) is crucial for understanding the nature of AGNs because it is believed that the mass accretion fuels a SMBH, supports its nuclear luminosity, and presumably powers its outflow. Given the simplest approximation of an axisymmetric adiabatic accretion flow onto a SMBH, Bondi (1952) derived a black hole accretion power $P_{\\rm B}=\\dot{M}_{\\rm B} c^2$, where $\\dot{M}_{\\rm B}$ is the Bondi accretion rate which is defined at the sphere of gravitational influence of the SMBH (i.e., the Bondi radius $r_{\\rm B}$).  The active galaxy M 87, one of the well-known LLAGNs, possesses a huge black hole mass $M_{\\bullet}=(3.2$--$6.6) \\times10^{9} M_{\\odot}$ (Macchetto et al. 1997; Gebhardt et al. 2011; Walsh et al. 2013) with a prominent relativistic jet\\footnote{We adopt 6.6 $\\times$ 10$^{9} M_{\\odot}$ for the black hole mass in this paper.}. The estimated jet power is in a range between $\\sim10^{42}$ to $10^{44}$ erg s$^{-1}$ (Li et al. 2009 and references therein) and seems to match fairly well with the overall energetics described by the Bondi accretion\\footnote{$\\dot{M}_{\\rm B} \\simeq 0.12 M_{\\odot}\\,$yr$^{-1}$ is adopted here.} ($P_{\\rm B} \\sim 7 \\times 10^{45}$ erg s$^{-1}$) at $r_{\\rm B}\\simeq 230$ pc. However, the central core of M 87 is highly under-luminous in the X-ray band and its X-ray luminosity is smaller than $P_{\\rm B}$ by five orders of magnitude ($L_{\\rm X} \\sim 7 \\times 10^{40}$ erg~s$^{-1}$; Di Matteo et al. 2003). The low X-ray luminosity implies the existence of radiatively inefficient accretion flows (RIAFs) and/or a substantial decrease of the mass accretion rate $\\dot{M}  \\ll \\dot{M_{\\rm B}}$ at radii $r \\ll r_{\\rm B}$.  The theory of RIAFs, such as the advection-dominated accretion flow (ADAF: Ichimaru 1977; Narayan \\& Yi 1995), convection-dominated accretion flow (CDAF: Narayan 2000; Quataert \\& Gruzinov 2000), and adiabatic inflow-outflow solution (ADIOS: Blandford \\& Begelman 1999) was developed over the last decades. The resulting mass accretion rates toward an SMBH normalized by the Bondi accretion rate can be scaled as a function of the spherical radius as $\\dot{M} /\\dot{M}_{\\rm B}=(r/r_{\\rm B})^{\\kappa}$ with $\\kappa=0 - 1$. Considering the viscosity parameter $\\alpha \\geq 0.01$ (Shakura \\& Sunyaev 1973), solutions of classical ADAFs suggest a mass accretion rate from $r_{\\rm B}$ to $r_{\\rm s}$ that is comparable to $\\dot{M}_{\\rm B}$ : $\\dot{M}_{\\rm GADAF} \\sim (0.1 - 1)\\times \\dot{M}_{\\rm B}$ (GADAF stands for ``giant'' ADAF; Narayan \\& Fabian 2011). On the other hand, CDAFs exhibit a substantial decrease of the mass accretion rate toward $r \\simeq r_{\\rm s}$ as $\\dot{M}_{\\rm CDAF} \\sim (r/r_{\\rm B}) \\times \\dot{M}_{\\rm B}$ (Igumenshchev \\& Narayan 2002). The ADIOS, which generalizes the ADAF model by including the disk wind, can take intermediate values of $\\kappa$ anywhere between 0 (ADAF) and 1 (CDAF). Recent numerical simulations favor $\\kappa \\sim 0.4$--0.7 which is consistent with ADIOS (e.g., Pang et al. 2011; Yuan et al. 2012).  There has been a growing consensus during the last decade that millimeter (mm)/submillimeter (submm) polarimetry provides useful diagnostics to infer $\\dot{M} $ of RIAFs at $r \\ll r_{\\rm B}$ (Agol 2000; Quataert \\& Gruzinov 2000). In particular, through the Faraday rotation measure (RM; an integral of the product of the thermal electron density and the magnetic field component along the line of sight) of the linear polarization, $\\dot{M_{\\bullet}}$ toward Sgr A* was examined (e.g., Bower et al. 2003; Marrone et al. 2006; Macquart et al. 2006) and an upper limit of $\\dot{M}$ was found to be $\\leq 10^{-7}$--$10^{-6} M_{\\odot}\\,$yr$^{-1}$ at $r \\lesssim 100\\, r_{\\rm s}$ ($r_{\\rm s}$ is the Schwarzschild radius).  We note that the submm wavelengths are more advantageous than other radio bands for determining the RM because the bandwidth depolarization is less significant, and the opacities of the accretion flow and the jet are smaller so that one can probe regions closer to the BH. However, efforts to determine the $\\dot{M} $ of AGNs (including M 87) have never been undertaken except for Sgr A*.   For M 87, all previous efforts on measuring the RM primarily focused on either the jet or radio lobes at centimeter wavelengths.  Owen et al. (1990) found RM values of a few thousands rad/m$^{2}$ at lobes (interpreted as arising from a foreground Faraday screen) and of a few hundreds rad/m$^{2}$ in the kiloparsec-scale jet.  More recently, Algaba et al. (2013) found similar RM values with a hint of gradients across the kpc jet.  On upstream parsec scales, Junor et al. (2001) found an average RM of $\\sim$$-$4400 rad/m$^{2}$ with regions that show both positive and negative signs. Zavala \\& Taylor (2002) found similar results and explained the origin of the observed RM with foreground medium not directly associated with the jet. To our knowledge, no mm/submm polarization observations for examining RM toward M 87 have been conducted.   In this letter, we present the first measurement of RM toward the M 87 nucleus at submm wavelengths. We also derive the mass accretion rate $\\dot{M}$ onto the black hole based on the measured RM and provide a constraint on RIAF models. In section 2, we introduce our SMA observation of the RM together with our data reduction. In section 3, we show our main results. The estimation of $\\dot{M}$, the discussion on the preferred RIAF models, and alternative interpretations of RM are presented in section 4. ",
        "conclusions": "\\begin{figure}[ht] \\begin{center} \\vspace*{0.0 cm} \\hspace*{0.0 cm} \\includegraphics[angle=0, scale=0.4]{fig4.ps} \\vspace*{0.0 cm} \\caption{The envisioned accretion flow in M 87 for deriving the rotation measure. We assume that the M 87 jet offers the background polarized emission for Faraday rotation, and the dominant contribution comes from the innermost core of the jet. $r_{in}$ and $r_{out}$ are the inner and outer edge of the Faraday screen in the accretion flow where the electrons are sub-relativistic and the magnetic field is coherent.} \\end{center} \\end{figure}  \\subsection{Constraints on Mass Accretion Rate and its Accretion Flow Model}  To constrain $\\dot{M}$ with RM, we essentially follow the procedure outlined in Marrone et al. (2006). The primary assumption in this method is that the hot accretion flow in front of a source of synchrotron emission causes Faraday Rotation. The model also assumes that the density profile of the accretion flow follows a power-law distribution ($n$ $\\propto$ $r^{-\\beta}$) and that the magnetic field is well ordered, radial and of equipartition strength.  In the case of M87, we further assume that the \\emph{innermost} jet (i.e. its base) provides the dominant background polarized emission in the submillimeter band (see Figure 4). Note that while polarized emission could be emitted along the jet all the way from close to the BH to radii beyond $r_{\\rm B}$, the dominant polarized emission at 230 GHz most likely originates from the jet base. This is because the majority of our SMA flux density (1.6 Jy) can be explained by the emission from the jet base at the center of M87. Based on the VLBI observation of M87 at 230 GHz, Doeleman et al. (2012) show that the jet base has a flux density of 1 Jy and a size of 5.5$r_{s}$. By using the core-shift relationship (Hada et al. 2011), Nakamura \\& Asada (2013) determine the offset of the jet base from the BH at 230 GHz to be 4.3$r_{s}$ (in projected distance). Therefore, these studies imply that the dominant background emission quite likely originates from within a few $r_{s}$ from the BH. Nonetheless, we currently cannot exclude the possibility of polarized emissions from the diffuse part of the jet, and this issue will be addressed with future high-resolution observations.  Based on Equation (9) in Marrone et al. (2006), we express $\\dot{M}$ as a function of RM as \\begin{equation*} \\dot M =2.2\\times10^{-9}\\left[1-(r_{out}/r_{in})^{-(3\\beta-1)/2}\\right]^{-2/3} \\times \\left( \\frac{M_{\\rm BH}}{6.6\\times 10^9 M_{\\odot}}\\right)^{4/3} \\left( \\frac{2}{3\\beta-1}\\right)^{-2/3} r_{in}^{7/6} RM^{2/3} ~\\rm rad/m^{2}, \\end{equation*} where $\\beta$ is a parameter depending on the accretion flow models ($\\beta$=1/2 $-$ 3/2 accounts for all varieties of accretion models). In the above equation, $r_{\\rm in}$ and $r_{\\rm out}$ are the inner and outer edge of the Faraday screen (in units of $r_{\\rm s}$) where the electrons are sub-relativistic and the magnetic field is coherent. We remark that the electrons within $r_{\\rm in}$ make a small or negligible contribution to RM because they become relativistic (i.e. $T_{e}>T_{\\rm Rel}=6\\times10^{9}$ K=$m_{e}c^{2}/k_{B}$) and RM is, thus, suppressed (Quataert \\& Gruzinov 2000).  In order to derive $\\dot{M}$ with the above equation, it is important to determine $r_{\\rm in}$ and $r_{\\rm out}$. We estimate $r_{\\rm in}$ based on the observed flux density of M87 from VLBI observations. As we described above, thermal electrons within $r_{\\rm in}$ become relativistic and can emit significant synchrotron radiation that dominates the emission from the accretion disk (Yuan et al. 2003). Since it is expected that a hot accretion flow would be optically thick at cm wavelengths, the size of the region with relativistic plasma can be constrained by using the observed flux density and brightness temperature. For M 87, the observed total flux density of the nucleus is 1 Jy from the VLBI observation at 43 GHz (Abramowski et al. 2012). Since the brightness temperature of the hot electrons within $r_{\\rm in}$ must be at least 6$\\times$10$^{9}$ K, the diameter (i.e. 2~$r_{\\rm in}$) of the emitting region of relativistic plasma can be constrained to be $\\leq$ 0.32 mas (= 42$r_{\\rm s}$). Since the observed flux density sets an upper limit on the amount of emission from relativistic electrons, it leads to an upper bound on $r_{\\rm in}$ of 21$r_{\\rm s}$. On the other hand, if the innermost jet resides right in front of the the emitting region of the relativistic plasma and the emission from the hot electrons is totally absorbed by the optically thick jet, we can not constrain the size of the region of the relativistic plasma directly from the observed flux. However, we can use the size of the innermost jet to constrain $r_{\\rm in}$.  The observed size of the VLBI core at 43 GHz is 17 $\\pm$ 4r$_{\\rm s}$ (Asada \\& Nakamura 2012). If the emission from the hot electrons is totally absorbed, $r_{\\rm in}$ must be smaller than 17 $\\pm$ 4r$_{\\rm s}$. Therefore, the 21$r_{\\rm s}$ upper bound for $r_{\\rm in}$ is reasonable even for this case. We fix $r_{\\rm out}$ to be at the Bondi radius $r_{\\rm B}$. Here, $r_{\\rm B}$ is just a convenient value to choose while $r_{\\rm out}$ is not well determined. This is a reasonable choice because $\\dot{M}$ is insensitive to $r_{\\rm out}$, and $\\dot{M}$ will only change by a factor of order unity even if $r_{\\rm out}$ $<<$ $r_{\\rm B}$. With these values, our new measurement of RM sets an upper limit $\\dot{M}$ of 9.2 $\\times$10$^{-4}$ M$_{\\odot}$~yr$^{-1}$ at 21$r_{\\rm s}$ for $\\beta \\leq 3/2$. This corresponds to 7.4 $\\times$ 10$^{-3}$ $\\dot{M}_{\\rm B}$, which is at least two orders of magnitude smaller than the Bondi accretion rate. This value implies that the mass accretion rate is significantly suppressed while material is accreted and falling in from $r_{\\rm B}$ to 21$r_{\\rm s}$.  Determining $\\dot{M}$ provides an effective way to constrain accretion flow models because each model has its own unique $\\dot{M}$ profile. The ADAF model requires $\\dot{M}$ to be comparable to $\\dot{M}_{\\rm B}$ from $r_{\\rm B}$ all the way to $r$$\\sim$$r_{\\rm s}$, while the ADIOS and the CDAF models require suppression of $\\dot{M}$ as ($r/r_{\\rm B}$)$^{\\kappa}$$\\dot{M}_{\\rm B}$, where $r$ is the radius from the BH with 0 $<$ $\\kappa$ $<$ 1 for ADIOS and $\\kappa=1$ for CDAF. Since our new results suggest significant suppression of $\\dot{M}$ at the radius of 21$r_{\\rm s}$, they disfavor the possibility of a classical ADAF and prefer ADIOS/CDAF. This result is in good agreement with the findings from nearly all numerical simulations (Yuan et al. 2012).  The maximum radiative power that can be extracted from mass accretion is $P_{rad}$ $=$ $\\eta \\dot{M} / c^{2}$ $\\simeq$ 7 $\\times$10$^{45}$ $\\eta$ $\\dot{M} / \\dot{M}_{\\rm B}$ erg~s$^{-1}$, where $\\eta$ is the radiative efficiency. With our upper limit of 9.2 $\\times$10$^{-4}$ M$_{\\odot}$~yr$^{-1}$ (= 7.4 $\\times$ 10$^{-3}$ $\\dot{M}_{\\rm B}$), the maximum radiative power is estimated to be 5 $\\times$ 10$^{42}$ erg~s$^{-1}$ if $\\eta=0.1$. This is about two orders of magnitude larger than the observed X-ray luminosity $L_{\\rm X}$ of 7$\\times$10$^{40}$ erg~s$^{-1}$, suggesting that the current limit of mass accretion rate is sufficient to explain the observed X-ray luminosity. On the other hand, the X-ray luminosity is low compared to the maximum available nuclear power. If $\\dot{M}$ is not significantly lower than the upper limit obtained here, this will imply that the radiative efficiency is significantly lower than the canonical value of 0.1. In this case, both mass accretion rate and radiative efficiency are significantly suppressed which is consistent with the expected properties of RIAFs.  \\subsection{Alternative interpretations ?} Our upper limit on the rotation measure is derived from a set of polarization position angles (PA) with still significant errors. Because of the large uncertainties, it is difficult to demonstrate that PAs follow the $\\lambda^{2}$ law, which is the basis of our fundamental assumption on the location of the Faraday screen (relative to the background source). As a result, we cannot completely rule out the scenario that $\\emph{internal}$ Faraday rotation (i.e. Faraday screen intermixed with emitting plasma in the jet) plays an important role and the possibility that we are measuring an RM purely originating from cold electrons in the sheath of the jet, especially because the jet in M87 points to the observer nearly along the line of sight (i.e. $<$ 20$^{\\circ}$).  To assess the possibility of a jet-based RM, we adopt the Burn model (Burn 1966) on internal Faraday rotation and depolarization. By using the degree of linear polarization of the M87 core from the current study (1 \\%) and that from the optical measurements (1\\%-13\\%; Perlman et al. 2011), we find that the maximum possible observed RM for the jet at 230 GHz is $\\sim$6$\\times$10$^{6}$ rad/m$^{2}$, beyond our current upper limit. Therefore, a jet-originated RM could provide an explanation for our observed RM.  Nonetheless, in the context of our simplified one-dimensional model for spherically symmetric RIAFs, we still prefer the scenario that the Faraday rotating plasma (i.e. accreting gas) is external to the emission region because the dominant polarized emission quite likely comes from within a few $r_{s}$ from the BH (see section 4.1). In this case, since all the non-relativistic electrons reside beyond $\\sim$21$r_{s}$ from the black hole, the Faraday screen must be external to the background source.  In essence, in order to constrain $\\dot{M}$ more unambiguously with the method used in this Letter, it is important to establish the $\\lambda^{2}$ law of PA and obtain a more precise RM. This will rely on higher sensitivity observations with a significantly longer lever arm in the frequency space and higher angular resolution in the future."
    },
    "1402/1402.6196_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small  The cycling of material from the interstellar medium (ISM) into stars and the return of stellar ejecta into the ISM is the engine that drives the {\\it galactic ecology} in normal spirals.  This ecology is a cornerstone in the formation and evolution of galaxies through cosmic time. There remain major observational and theoretical challenges in determining the processes responsible for converting the low-density, diffuse components of the ISM into dense molecular clouds, forming  dense filaments and clumps, fragmenting them into stars, expanding OB associations, and bound clusters, and characterizing the feedback that limits the rate and efficiency of star formation. This formidable task can be attacked effectively for the first time thanks to the synergistic combination of new global-scale surveys of the Milky Way from infrared to radio wavelengths, offering the possibility of bridging the gap between local and extragalactic star formation studies.  The Herschel Hi-GAL survey, with its five-band 70--500\\mum\\ full Galactic Plane mapping at 6--36\\arcsec\\ resolution, is the keystone of a set of continuum surveys that include GLIMPSE(360)+MIPSGAL@Spitzer, WISE and MSX, ATLASGAL@APEX, BGPS@CSO, and CORNISH@VLA. This suite enables us to measure the Galactic distribution and physical properties of dust on all scales and in all components of the ISM from diffuse clouds to filamentary complexes and hundreds of thousands of dense clumps. A complementary suite of spectroscopic surveys in various atomic and molecular tracers is providing the chemical fingerprinting of dense clumps and filaments, as well as essential kinematic information to derive distances and thus transform panoramic data into a 3D representation.  The latest results emerging from these Galaxy-scale surveys are reviewed.  New insights into cloud formation and evolution, filaments and their relationship to channeling gas onto gravitationally-bound clumps, the properties of these clumps, density thresholds for gravitational collapse, and star and cluster formation rates are discussed. \\\\~\\\\~\\\\~}   % ",
        "introduction": "Phase changes in the Galactic ISM are to a large extent controlled by the formation of massive stars.  The cycling of the ISM from mostly neutral atomic (\\HI) clouds into molecular (H$_2$) clouds, traced by low-excitation species such as OH and CO, leads to the formation of dense, self-gravitating clumps and cores traced by high-density species such as NH$_3$, CS, HCN, HCO$^+$, and N$_2$H$^+$, and by other even higher dipole moment molecules in regions where stars form.  Dust continuum emission in the mid-IR, far-IR, sub-millimeter (submm), and millimeter ranges of the spectrum reveal progressively cooler and higher column density dust associated with all phases of the ISM, and with the high-density material in protostellar envelopes and disks. Dust extinction measurements in the UV, visual, near-IR, and mid-IR enable detection of progressively higher column densities of dust located in front of background stellar and diffuse emission sources.  Feedback from young stars limits the rate and efficiency of star formation by generating turbulence and disrupting the parent clouds. In giant molecular clouds (GMCs) feedback, usually dominated by the most massive young stellar objects (YSOs), tends to disrupt the parent cloud by the time 2--20\\%\\ of its mass has formed stars.  Ionization and shocks produced by massive stars, associations, and clusters convert the remaining gas into the $10^4$--$10^8$~K photoionized and shock-heated phases of the ISM -- \\HII\\ regions traced by H and He recombination lines and free-free emission, hot superbubbles traced by X-ray emission, and 0.1 to 1 kpc-scale supershells -- that eventually cool, condense, and reform the cool $\\sim$ $20$ to $10^4$~K \\HI\\ phase.  Compression by shocks and gravity leads to the formation of new GMCs.  In the Solar vicinity, atoms cycle through this loop on a time scale of 50--100~My.  A mean star formation efficiency (SFE) of 5\\% per GMC implies that atoms on average pass through this loop about 20 times before being incorporated into a star.  This {\\it galactic ecology} is modulated by large-scale processes such as spiral arms and the central bar of the Milky Way.  While star formation depletes the ISM at a rate of $\\sim$ 2 \\msol~y$^{-1}$, infall of gas from the Local Group supplements it at a highly uncertain rate of 0.1 to 1 \\msol~y$^{-1}$. The balance between star formation, which sequesters matter for the main-sequence lifetime of stars, and recycling from stellar winds and dying stars, supplemented by infall from the Local Group, determines the time scale on which the Galactic ISM is depleted.  Large-scale Galactic plane surveys of the last decade provide the data required to flesh out the details of this {\\it galactic ecology}.  The Milky Way, a mildly-barred, gas-rich spiral galaxy, supports the most active star formation in the Local Group.  The Galactic disk contains about 1--$3 \\times 10^9$ M$_{\\odot}$ of H$_2$ and about 2--$6 \\times 10^9$ M$_{\\odot}$ of \\HI\\ \\citep{Combes1991,KalberlaKerp2009}. The molecular disk becomes prominent from Galactocentric radius $R_{gal} \\sim$ 3 kpc, beyond the central bar, with surface density peaking at about $R_{gal} \\sim$ 4--6 kpc (often referred to as the {\\it Molecular Ring}).  It declines exponentially toward larger radii but can be traced well beyond $R_{gal} \\sim $10 kpc. The \\HI\\ disk extends to beyond $R_{gal} \\sim$ 20 kpc. The disk ISM is dominated by a four-arm spiral pattern consisting of two major and two minor spiral arms and a variety of inter-arm spurs. The Sun is currently located in an inter-arm region between the Sagittarius and Perseus spiral arms, near a spur which extends from a distance of several kpc in the direction of Cygnus to at least 1 kpc beyond Orion in the opposite direction \\citep{Xu2013}. In molecular tracers such as CO, the arm--inter-arm contrast in the Molecular Ring is around 3:1, but in the outer Galaxy beyond the Solar circle at $R_{gal} \\approx$ 8.5 kpc the contrast is much higher, approaching a value of 40:1 toward the Perseus arm.  Although Galactic rotation is well described by circular motions with orbit speeds ranging from 200 to 250 km~s$^{-1}$ from $R_{gal} \\sim$ 1 kpc to beyond 20 kpc, substantial radial motions are seen in gas tracers toward the Galactic Center and Anti-center.  The line of sight toward the Galactic Center shows three clearly-defined spiral arms, distinguished by their negative radial velocities with respect to the local standard of rest (LSR; defined by the mean motion of stars in the Solar vicinity).  The innermost is the so-called 3-kpc arm in the Molecular Ring ($V_{LSR} \\approx -60$ km~s$^{-1}$), next is the Scutum arm ($V_{LSR} \\approx -35$ km~s$^{-1}$), and then closest ($\\sim$ 2 kpc) is the Sagittarius arm ($V_{LSR} \\approx -15$ km~s$^{-1}$). Radial streaming motions of order 10 to 30 km~s$^{-1}$ are also seen toward the Perseus arm and the far outer arm beyond that.  These motions might reflect the gravitational potential well depth of the spiral arms.  The inner edge of the Molecular Ring at $R_{gal} \\sim 3$~kpc lies near the outer radius of the Galactic bulge and bar.  The central 3-kpc region of the Galaxy is dominated by the stellar bar whose major axis is inclined by 20\\deg\\ to 40\\deg\\ with respect to our line of sight \\citep{binney91,MorrisSerabyn1996}.  Within 0.5 kpc of the nucleus lies the Central Molecular Zone (CMZ) containing $\\sim 10$\\% of the molecular gas in the Galaxy.  The CMZ clouds are one to two orders of magnitude denser than GMCs in the Molecular Ring, and an order of magnitude more turbulent. ",
        "conclusions": ""
    },
    "1402/1402.1549_arXiv.txt": {
        "abstract": "We have investigated the interstellar reddening law of young open clusters within 3kpc from the Sun using optical, near-IR 2MASS, and {\\it Spitzer} IRAC data. The total-to-selective extinction ratio $R_V$ of 162 young open clusters ($\\log ~t_{age} \\lesssim 7.3$) listed in the open cluster database WEBDA is determined from the color excess ratios. The young open clusters in the Sgr-Car arm show a relatively higher $R_V$, those in the Per arm and in the Cygnus region of the local arm show a relatively smaller value, and those in the Mon-CMa region of the local arm show a normal value ($R_V \\approx 3.1$). ",
        "introduction": "To determine the physical parameters of stars and clusters accurately, we should know the correct interstellar reddening. Precise knowledge of $R_V$ is of crucial importance for the determination of distances of reddened stars or clusters. Many investigators determine the total extinction $A_V$ by assuming the total-to-selective extinction ratio $R_V$ to be 3.1. In 1940s, from multicolor photometry of many O and B type stars \\citet{sw43} concluded that the law of selective absorption is the same for all directions in the Galaxy. Later \\citet{jb63} conducted multi-color photometry of O, B stars and showed that there is no unique reddening law. In addition, they showed that there is a minor variation of $R_V$ with galactic longitude. Later \\citet{serkowski75} confirmed such a variation from the spatial variation of wavelength of maximum polarization. \\citet{whittet77} presented a functional form of $R_V$ variation against galactic longitude. Recently \\citet{fitzpatrick09} showed that the interstellar reddening law $R_V$ is different for different sightlines.  Recently \\citet{hur12} showed that an abnormal reddenig law for the intracluster medium of the young open clusters in the $\\eta$ Carina nebula. However they obtained a fairly normal reddening law for the foreground stars. On the other hand, \\citet{sung13a} found a nearly normal reddening law toward the young open cluster NGC 6231 from the optical to mid-IR photometric data. Now it is possible to study the reddening law toward young open clusters in a homogeneous way because the homogeneous near-IR 2MASS data as well as mid-IR {\\it Spitzer} IRAC images are available for many young open clusters. In this paper we present a preliminary result from our systematic investigation of the reddening law.  \\begin{figure*}[!t] \\begin{center} \\plotone{sungh_f1.ps} \\caption{Reddening law of Trumpler 16 from the optical to mid-IR. The reddening for a given color with respect to $E(B-V)$ is shown. The optical and near-IR data are from \\citet{hur12} and 2MASS, respectively, while mid-IR photometry was obtained from the {\\it Spitzer} IRAC images. The solid and dashed lines represent the normal reddening law for the foreground ($R_{V,fg} = 3.1$) and the abnormal reddening law for the intracluster medium of Trumpler 16 region ($R_{V,cl} = 4.6$), respectively. Large dots and squares denote stars with higher proper motion membership probability ($P_\\mu \\geq 70$) and with lower probability ($P_\\mu < 30$), respectively. \\label{fig1} } \\end{center} \\end{figure*} ",
        "conclusions": "From the analysis of photometric data for 162 young open clusters we confirmed the large scale variation of the reddening law in the galactic plane, which implies the real variation of dust size distribution in the galactic plane. We also found a large fluctuation of $R_V$ for a given region. The cause of this fluctuation is still uncertain, and therefore a systematic photometric survey of open clusters to provide homogeneous photometric data, e.g. the Sejong Open cluster Survey (SOS) \\citep{sung13b} is required to confirm the real fluctuation.  \\begin{figure*}[!t] \\begin{center} \\plotone{sungh_f2.ps} \\caption{Plot of $R_V$ against galactic longitude $l$. The solid curve is the variation of $R_V$ from \\citet{whittet77}. Large dots represent the $R_V$ determined from the color excess ratios, while open circles denote the $R_V$ determined from less reliable data. The size of symbols is proportional to the number of OB stars used in the determination of $R_V$.  \\label{fig2} } \\end{center} \\end{figure*}"
    },
    "1402/1402.1852_arXiv.txt": {
        "abstract": "The quantitative measurement of O\\emissiontype{VII} line intensity is a powerful method for understanding the soft X-ray diffuse background. By systematically analyzing the O\\emissiontype{VII} line intensity in 145 high-latitude Suzaku/XIS observations, the flux of O\\emissiontype{I} fluorescent line in the XIS spectrum, contaminating the O\\emissiontype{VII} line, is found to have an increasing trend with time especially after 2011. For these observations, the O\\emissiontype{VII} line intensity would be overestimated unless taking into consideration the O\\emissiontype{I} fluorescent line contamination. Since the O\\emissiontype{I} line emission originates from solar X-rays, this increase suggests that incident solar X-rays at the O\\emissiontype{I} fluorescence energy tend to be larger than the early phase of Suzaku observations (2005 -- 2010). ",
        "introduction": "The soft X-ray diffuse background (SXDB; \\cite{1990ARA&A..28..657M}) is distributed over the whole sky. In the SXDB below 1 keV, the blend of the O\\emissiontype{VII} K$\\alpha$ triplet line emission (centroids of the resonance, inter-combination and forbidden lines: 0.574, 0.569 and 0.561 keV, respectively) is one of the most prominent constituents \\citep{2002ApJ...576..188M}. \\citet{2009PASJ...61..805Y} suggests that it arises from two origins: approximately uniform emission of $\\sim2$ photons s$^{-1}$ cm$^{-2}$ sr$^{-1}$ (LU), and spatially variable emission (0 -- 7 LU) from hot gas of temperature $kT\\sim0.2$ keV. The former is considered to originate from the heliospheric solar wind charge exchange (SWCX; \\cite{1998LNP...506..121C, 2000ApJ...532L.153C, 2004A&A...418..143L}) plus the local hot bubble (LHB; \\cite{1990ARA&A..28..657M}), and the latter from hot gas of the Galactic halo \\citep{2000ApJ...543..195K}. \\citet{2011ApJ...726...91K} provides constraints on the heliospheric SWCX$+$LHB contributions to the local SXDB. \\citet{2013PASJ...65...32Y} gives an indication of long-term time variability of the O\\emissiontype{VII} line intensity in the heliospheric SWCX due to the 11-year solar activity. The physical state of the hot gas of the Galactic halo has been investigated with O\\emissiontype{VII} absorption lines (e.g. \\cite{2008ApJ...672L..21Y}; \\cite{2010PASJ...62..723H}; \\cite{2012ApJ...756L...8G}; \\cite{2013ApJ...770..118M}). The O\\emissiontype{VII} line is used as a probe to understand about surrounding plasma around the solar system and the Galaxy.  For the study of the O\\emissiontype{VII} line emission in the SXDB, X-ray Imaging Spectrometers (XISs; \\cite{2007PASJ...59S..23K}) aboard Suzaku \\citep{2007PASJ...59S...1M} are best suited because of their low and stable non X-ray background (NXB) and better energy response than those of CCDs onboard XMM-Newton or Chandra. The O\\emissiontype{VII} line can be contaminated by the O\\emissiontype{I} K$\\alpha$ line (centroid: 0.525 keV) because CCDs such as Suzaku/XIS cannot completely distinguish the O\\emissiontype{I} line from the O\\emissiontype{VII} line owing to their energy resolution of $\\sim0.1$ keV (FWHM at $\\sim1$ keV). The O\\emissiontype{I} line is created by fluorescence of solar X-rays with neutral oxygen in the Earth's atmosphere. It is prescriptively thought to be removable by discarding time periods when the telescope is looking at the bright Earth or its atmosphere (e.g. elevation angle from bright Earth limb $>20^\\circ$).  We analyzed 145 observations with the Suzaku/XIS from 2005 to 2012 and investigated the O\\emissiontype{VII} line intensity in the SXDB with careful attention to the O\\emissiontype{I} fluorescent line contamination. In this letter, we report on an increasing trend of the O\\emissiontype{I} fluorescent line intensity with time, especially after 2011, found during the analysis. The error ranges stated in this letter show 90\\% confidence level from the center value, and the vertical error bars in the figures indicate 1$\\sigma$ level. ",
        "conclusions": "The O\\emissiontype{I} line increased as significantly as detectable in Suzaku/XIS observations of the SXDB after 2011. The origin of it is considered to be the Earth's atmosphere because of dependence on DYE\\_ELV and correlation between the O column density and the 0.48 -- 0.55 keV count rate. The best way to identify possible contamination and to remove the O\\emissiontype{I} emission is to fold the light curve by the orbital period, or to sort by the O column density. Adding a Gaussian for the O\\emissiontype{I} in the spectral fitting may help exclude possible O\\emissiontype{I} contamination when evaluating the O\\emissiontype{VII} intensity.  The O\\emissiontype{I} intensity ($I_{\\rm O\\emissiontype{I}}$) is following equation: \\begin{equation} I_{\\rm O\\emissiontype{I}}=\\frac{1}{4\\pi}f_{\\rm O\\emissiontype{I}}N_{\\rm O}\\int F_{\\rm Sun}(E)\\sigma_{\\rm O\\emissiontype{I}}(E)dE \\end{equation} where $f_{\\rm O\\emissiontype{I}}$, $N_{\\rm O}$, $F_{\\rm Sun}(E)$ and $\\sigma_{\\rm O\\emissiontype{I}}(E)$ are the fluorescence yield of O\\emissiontype{I} (0.0058; \\cite{1979JPCRD...8..307K}), the O column density, the flux of the Sun at the Earth distance specified as a function of photon energy and the cross section for inner-shell ionization of O, respectively. $F_{\\rm Sun}(E)$ is derived from \\citet{2001ApJ...560..499O} and $\\sigma_{\\rm O\\emissiontype{I}}(E)$ is described in the Evaluated Nuclear Data File\\footnote[3]{\\url{http://t2.lanl.gov/nis/data/endf/endfvii-atomic.html}}. Solar X-rays are mainly emitted from quiet regions and active regions of the solar corona. The quiet regions produce nearly steady activities through the solar minimum and maximum. Considering the contribution of the quiet regions, the energy integral of $F_{\\rm Sun}(E)\\sigma_{\\rm O\\emissiontype{I}}(E)$ is about $2\\times10^{-12}$ photons s$^{-1}$, and $I_{\\rm O\\emissiontype{I}}$ is \\begin{equation} I_{\\rm O\\emissiontype{I}}\\sim10^{-15}\\ {\\rm photons\\ s^{-1}\\ sr^{-1}}\\times N_{\\rm O}. \\end{equation} Thus, $N_{\\rm O}=10^{15}$ cm$^{-2}$ implies $I_{\\rm O\\emissiontype{I}}\\sim1$ LU. This calculation result supports the fact that the O\\emissiontype{I} fluorescent line is seen when $N_{\\rm O}>10^{15}$ cm$^{-2}$ in the early phase of Suzaku observations (2005 -- 2010); under the solar minimum. On the other hand, the active regions of the solar corona significantly appear during the solar maximum. The flux from the active regions is more than tenfold larger than that from the quiet regions. It is compatible with the fact that the O\\emissiontype{I} clearly appears when $N_{\\rm O}>10^{14}$ cm$^{-2}$ in some observations after 2011 (e.g. LOCK 365); under the solar maximum. In the other observations after 2011, no clear the O\\emissiontype{I} fluorescent line is seen even when $N_{\\rm O}>10^{14}$ cm$^{-2}$. The reason can be that position and property of the active regions of the solar corona differ day by day. Based on the above, a possible explanations for increasing tendency of the O\\emissiontype{I} fluorescent line is conceivable; incident solar X-rays tend to be larger than the early phase of Suzaku observations (2005 -- 2010).  \\bigskip We are grateful to Dr. Randall K. Smith, the referee, for fruitful suggestions. We acknowledge Kazuhiro Sakai for helpful comments about the analysis. We also thank to the Suzaku, ACE and WIND team for providing the data. This work is partly supported by a Grants-in-Aid for Scientific Research from JSPS (Project Number: 12J10673)."
    },
    "1402/1402.3316_arXiv.txt": {
        "abstract": "We consider the single field chaotic $m^2\\f^2$ inflationary model with a period of preheating, where the inflaton decays to another scalar field $\\chi$ in the parametric resonance regime. In a recent work, one of us has shown that the $\\chi$ modes circulating in the loops during preheating notably modify the $\\lf\\zeta\\zeta\\rg$ correlation function. We first rederive this result using a different gauge condition hence reconfirm that superhorizon $\\zeta$ modes are affected by the loops in preheating. Further, we examine how $\\chi$ loops give rise to non-gaussianity and affect the tensor perturbations. For that, all cubic and some higher order interactions involving two $\\chi$ fields are determined and their contribution to the non-gaussianity parameter $f_{NL}$ and the tensor power spectrum are calculated at one loop. Our estimates for these corrections show that while a large amount of non-gaussianity can be produced during reheating,  the tensor power spectrum receive moderate corrections. We observe that the loop quantum effects increase with more $\\chi$ fields circulating in the loops indicating that  the perturbation theory might be  broken down. These findings demonstrate that the loop corrections during reheating are significant and they must be taken into account for precision inflationary cosmology. ",
        "introduction": "Single scalar field inflationary models have solid predictions for the scalar and the tensor power spectra, and  the amount of non-gaussianity produced by the interactions. These observable quantities are fixed by a few parameters like the slow-roll parameter of the potential. Moreover, in these models the quantum loop corrections to the standard inflationary predictions turn out to be quite small (see  e.g. \\cite{s1,s2,s3}). As a result of this firm structure, many single field scalar models are either ruled out or severely constrained by the recent Planck data \\cite{pl} (see \\cite{ls} for a scan of inflationary scenarios in the light of Planck). For example, the chaotic $m^2\\f^2$ model is ruled out by 95\\% confidence level (provided the index is not running) by the contours in the the scalar-to-tensor ratio $r$ vs. the scalar spectral index $n_s$ data plane \\cite{pl}.  The constancy of the superhorizon curvature perturbation $\\zeta$ is very crucial for the inflationary predictions to hold. This helps to determine the cosmic microwave background (CMB) fluctuations from the correlation functions evaluated at the horizon crossing time. Technically, the conservation of $\\zeta$ sets an upper bound for the time integrals that appear in the in-in perturbation theory \\cite{mal}.  On the other hand, it is well known that the entropy perturbations can cause superhorizon evolution of the curvature perturbation (see e.g. \\cite{mfb}). In reheating, those fields that are unimportant during the exponential expansion are exited by the inflaton decay and they start to dominate the universe. Moreover, in some models the decay can occur violently in a preheating stage \\cite{reh1,reh2,reh3,reh4}. Although these entropy perturbations are not produced at cosmologically interesting scales, reheating stage ends with highly nonlinear processes  (see e.g. \\cite{n1,n2,n3}). While these nonlinearities can be effectively described by fluid dynamics that only affect local quantities \\cite{bsz}, some of them are known to have important consequences (see e.g. \\cite{coh1,coh2}).  Classical, and similarly quantum, nonlinearities imply that Fourier modes do not evolve independently. As a result of  this mode-mode coupling, short distance fluctuations are expected to affect the long wavelength modes. For example, a cubic interaction term in a Lagrangian would allow two modes with nearly equal large momenta to change the amplitude of a mode with small momentum. In quantum theory, there are also virtual modes circulating in the loops that affect the correlation functions. Evidently, it is crucial to determine the size of such effects. In a recent work, one of us has shown that in the chaotic $m^2\\f^2$ model with a period of preheating where the inlaton $\\f$ decays to another scalar $\\chi$, the parametrically amplified $\\chi$ modes appearing in the loops would meaningfully modify the curvature power spectrum \\cite{ali1}. This is an example of the entropy perturbations affecting the superhorizon curvature variable, however not by the real physical fluctuations but because of the virtual entropy modes appearing in the loops (entropy modes are known to give power  loop infrared divergences during inflation \\cite{ir1,ir2}).  Since $\\zeta$ becomes an ill defined dynamical variable during reheating, the calculations in \\cite{ali1} have been carried out in the $\\zeta=0$ gauge. In that case, one may first  calculate the $\\chi$ loop corrections to the inflaton fluctuation power spectrum until  the coherency of the  inflaton oscillations is lost.  After that moment, the possible effects on the superhorizon evolution are expected to be averaged out and become negligible. One may then apply a gauge transformation to read the $\\zeta$ power spectrum. Since in the first stage of the preheating the  background inflaton  oscillates coherently, the superhorizon modes are affected without violating causality \\cite{ca}.  Note that as long as the relativistic equations are treated properly, there should not arise any issue with causality.  In this paper, our first aim is to carry out the calculation of \\cite{ali1} in the $\\vf=0$ gauge, i.e. we will use $\\zeta$ directly as the main dynamical variable. Because $\\zeta$ is only ill defined at isolated times when the inflaton velocity vanishes, the propagator has ``spikes\" and it diverges at these moments.  We smooth out these spikes by using the time averaged background quantities in the $\\zeta$ action.  In \\cite{ali1}, only the loops arising from the interaction potential have been considered. Here, we determine all cubic interactions involving $\\chi$ and $\\zeta$, and estimate the total one loop correction to the $\\zeta$ power spectrum. Not surprisingly, our computations confirm the findings of \\cite{ali1} and show a significant contribution to the  $\\lf\\zeta\\zeta\\rg$  correlation function.  Our second aim in this work is to determine the  $\\chi$ loop contributions to the non-gaussianity and to the tensor power spectrum (in single field models, the bi-spectrum is not altered by the parametric resonance effects \\cite{yeni1}). Notable modifications to the scalar power spectrum found in \\cite{ali1} indicate the existence of similar significant corrections for these observables. We estimate the amount of non-gaussianity from the $\\zeta$-three point function by calculating the one loop graphs arising from the cubic interactions. It turns out that these corrections to the three point function can be expressed in terms of the two point function and it is possible to read the shape independent non-gaussianity order parameter $f_{NL}$. Similarly, the $\\chi$ field  coupled to the tensor fluctuations yield loop corrections to the tensor power spectrum. Since the tensor field behaves like a test field propagating on the background, the tensor calculation is not affected by different time slices of spacetime. We find that the tensor power spectrum is moderately corrected by the loops in reheating.  The organization of the paper is as follows. In the next section, we consider the chaotic $m^2\\f^2$ model with the extra preheating scalar field $\\chi$ to which the inflaton decays in the parametric resonance regime. In section \\ref{iii},  we  determine the cubic interaction terms involving the curvature perturbation $\\zeta$, the tensor mode $\\cc_{ij}$  and the preheating scalar $\\chi$.  We then calculate one loop corrections to the scalar power spectrum $\\lf\\zeta\\zeta\\rg$, the three-point function $\\lf\\zeta\\zeta\\zeta\\rg$, the tensor power spectrum $\\lf\\cc\\cc\\rg$ and make order of magnitude estimates of these corrections by using the theory of preheating. In section \\ref{sec3}, we further consider some higher order interactions involving two $\\chi$ fields and determine their loop effects. In \\ref{sec4} we conclude with  remarks and future directions. ",
        "conclusions": "\\label{sec4}  In a recent work \\cite{ali1}, one of us has shown that the loop quantum effects during reheating significantly modify the scalar power spectrum. In this paper, in an attempt to extend the findings of \\cite{ali1} we consider how loops in reheating produce non-gaussianity and affect the tensor power spectrum in the chaotic $m^2\\f^2$ model. Based on the tree level results, this model is actually ruled out by 95\\% confidence level by the Planck data (provided the running of the index is neglected), however our findings show that quantum effects during reheating can change this conclusion since the corresponding corrections can alter the tree level results appreciably.  In most of the scalar field inflationary models, inflation is followed by a period of coherent inflaton oscillations where the background is still homogeneous and isotropic. This phase continues until the backreaction effects are set in. As pointed out  in \\cite{ca}, in such a background causality does not preclude the emergence of  the superhorizon effects because by coherency the same physical influence can appear at different positions at the same time.  Therefore,  the quantum effects can be important for cosmological variables in the first stage of reheating. On the other hand, it is known that the entropy perturbations can cause nontrivial superhorizon evolution of the curvature perturbation. Consequently, it is not surprising to see that the effects of entropy modes circulating in the loops become significant, especially  in the parametric resonance regime. Indeed, we observe that  the corrections get larger as the number of $\\chi$ modes circulating in the loops increases, which indicates that the perturbation theory might become invalid.  It is well known that in the chaotic model we have studied, the curvature perturbation $\\z$ becomes an ill defined variable during reheating. Because of that reason in \\cite{ali1}, the calculations have been carried out in the $\\z=0$ gauge till the end of the first stage of reheating and then a gauge transformation has been applied to read the $\\lf\\z\\z\\rg$ correlation function. In this paper, we utilize a different strategy and smooth out the spikes of $\\z$ by using the averaged out background variables in the quadratic $\\z$-action. As it is shown above, the results obtained in this way is consistent with \\cite{ali1} and thus our conclusions about the scalar power spectrum (and non-gaussianity) are firm. Note that the tensor calculation is free from the gauge fixing issues.  It is possible to develop the results of this paper in different directions. Due to the importance of the chaotic $m^2\\f^2$ model,  it would be valuable to perform a full numerical check  of the loop corrections that are estimated in this paper. It would also be crucial to see whether the loops in reheating modify the predictions of the models that are favored by Planck data, like the Starobinsky model \\cite{strm}.  Finally, it would be interesting to determine the loop effects when the inflaton decay occurs perturbatively. In that case while the reheating scalar modes cannot take large values, the decay process is completed in a long time that might enhance the quantum effects, since according to in-in formalism \\eq{inp}, the quantum corrections are proportional to the duration of the process."
    },
    "1402/1402.1639_arXiv.txt": {
        "abstract": "\\noindent  We accurately determine the fundamental system parameters of the neutron--star X-ray  transient \\target\\ solely using phase-resolved high--resolution UVES spectroscopy. We first determine the radial-velocity curve of the secondary star and then model the shape of the phase-resolved absorption line profiles using an X-ray binary model. The model computes the exact rotationally broadened phase-resolved spectrum and does not depend on assumptions about the rotation profile, limb-darkening coefficients and the effects of contamination from an accretion disk. We determine the secondary star-to-neutron star binary mass ratio to be 0.1755$\\ppm$0.0025, which is an order of magnitude more accurate than previous estimates. We also constrain the inclination angle to be 32$^\\circ$$^{+8^\\circ}_{-2^\\circ}$. Combining these values with the results of the radial velocity study gives a neutron star mass of 1.94$^{+0.37}_{-0.85}$\\Msun \\ consistent with previous estimates. Finally, we perform the first Roche tomography reconstruction of the secondary star in an X-ray binary. The tomogram reveals surface inhomogeneities that are due to the presence of cool starspots. A large cool polar spot, similar to that seen in Doppler images of rapidly-rotating isolated stars is present on the Northern hemisphere of the K7 secondary star and we estimate that $\\sim$4 per cent of the total surface area of the donor star is covered with spots. This evidence for starspots supports the idea that magnetic braking plays an important role in the evolution of low-mass X-ray binaries. ",
        "introduction": "\\label{intro}  In interacting binaries, such as dwarf novae and X-ray transients, the secondary star's rotational broadening (\\vsini) combined with the radial velocity semi-amplitude ($K_{\\rm 2}$) is normally used to determine the binary mass ratio of the system $q$ (=$M_2$/$M_1$: where $M_{\\rm 1}$ and $M_{\\rm 2}$ are the mass of the compact and secondary star, respectively). However, due to the faintness of the secondary star, its rotational broadening is usually determined by using intermediate resolution ($\\sim$0.5--1.0\\,\\AA) spectroscopy, where \\vsini\\ is typically 30--100\\kms. With intermediate-resolution spectroscopy the information about the shape of the absorption lines is lost, and therefore the only information that can be extracted is the amount by which the stellar spectrum is broadened. The procedure commonly used to measure the secondary star's rotational broadening is to compare it to the spectrum of a slowly rotating template star, observed with the same instrumental configuration, that has been convolved with a limb-darkened standard rotation profile \\citep{Gray92}. One assumes a limb-darkening coefficient for the spectral line and usually  adopts zero limb-darkening or the continuum value (which depends on the wavelength and the star's effective temperature). The width of the standard rotation profile is varied until an optimum match is found with the target spectrum \\citep[see][]{Marsh94}.  \\citet{Shahbaz98} showed that, in principle, it is possible to use the shape of the secondary star's absorption lines to determine the binary system parameters. In particular, one can determine the binary mass ratio directly by comparing the secondary star's line profile with a model line profile for the Roche-lobe filling secondary star geometry. Indeed, \\citet{Shahbaz03} presented a model \\textsc{XrbSpectrum}, for determining the mass ratio of interacting binaries by directly fitting the observed line profile with synthetic spectra. The author makes direct use of \\textsc{NextGen} model atmosphere intensities (Hauschildt, Allard, \\& Baron 1999), which are the most comprehensive and detailed models available for cool stars. The model fully takes into account the varying temperature and gravity across the secondary star's photosphere, by incorporating the synthetic spectra into the secondary star's Roche geometry. As a result, \\citet{Shahbaz98} determine the exact rotationally broadened line profile of the secondary star and so eliminate the need for a limb-darkening law, and the uncertainties associated with it.   X-ray transients  are  a subset of  low-mass X-ray  binaries (LMXBs) that display episodic, dramatic X-ray  and optical  outbursts,  usually lasting for several months. Between outbursts, X-ray transients  remain in a quiescent state, with typical X-ray luminosities less than $10^{32}$\\erg, allowing the optical detection of the faint low-mass donor star.  This allows the possibility to perform radial velocity studies, probe the nature of the compact star and to also determine its mass  \\citep[e.g.][]{Charles06}.   The neutron star X-ray transient \\target\\ was discovered in 1969 by the Vela 5B satellite \\citep{Conner69} when it went into X-ray outburst. During its second outburst in 1979 the optical counterpart was discovered \\citep{Canizares80} and the mass accreting star was identified as a neutron star due to the fact that it displayed a Type I X-ray burst \\citep{Matsuoka80}. After 1980 \\target\\ stayed in quiescence at $V\\simeq 18.2$ and subsequent photometric studies led to the discovery of the orbital period $P_{\\rm orb}$\\,=\\,15.1\\,h \\citep{Chevalier89} and the determination of the companion star's radial velocity curve \\citep{Cowley88, MR90, Torres02, Avanzo05, Casares07}. Shahbaz, Naylor, \\& Charles (1993) modelled the quiescent $H$-band light curve as due to the Roche-lobe filling secondary star's ellipsoidal modulation and obtained a binary inclination in the range $i$=31--54$^\\circ$, assuming no contamination from the accretion disk. Using IR spectroscopy Khargharia, Froning, \\& Robinson (2010) determined the contribution of the secondary star to the infrared flux and remodelled the $H$-band light curve of  \\citet{Shahbaz93}, correcting for the fractional contribution of the donor star to obtain an inclination angle of 35$^\\circ$$^{+4^\\circ}_{-1^\\circ}$  \\citet{Torres02} obtained high-resolution optical spectroscopy of \\target\\ and  determined the mass function to be $f(M)$\\,=\\,0.220$\\ppm$0.005\\,\\Msun, which was later refined by \\citet{Avanzo05} to $f(M)$\\,=\\,0.201$\\ppm$0.004\\,\\Msun. \\citet{Avanzo06} searched for the effects of irradiation on the absorption-line radial-velocity curve and concluded that there is no evidence for irradiation. \\citet{Casares07} obtained much tighter constraints on $K_{\\rm 2}$\\,=\\,144.6$\\ppm$0.3\\kms\\ and $P_{\\rm orb}$\\,=\\,0.6290522$\\ppm$0.0000004\\,d and also determined \\vsini=44$\\pm$3\\,\\kms and $q$\\,=\\,0.20$\\ppm$0.03, using the standard method by comparing the observed spectrum with a template star convolved with a limb-darkened rotation profile. Estimates for the secondary star's spectral type using optical spectroscopy include a K7\\,V star \\citep{Shahbaz93} and a K3--K5\\,V star \\citep{Torres02, Avanzo05}. \\citet{Jonay05} determined $T_{\\rm eff}$\\,=\\,4500\\,K and  $\\log\\,g$\\,=\\,3.9, corresponding to a a K4\\,V star, as well as the chemical abundances.  The measurement of stellar masses in quiescent X-ray transients relies on the determination of $K_{\\rm 2}$ and $q$ from  spectroscopy    and $i$ normally from photometry (via ellipsoidal modulation studies of the  optical/IR lightcurves). This is mainly due to the fact that the secondary stars  in quiescent X-ray transients are optically extremely faint, making it  difficult for one to determine  accurate masses from spectroscopy alone. In non-eclipsing systems the main uncertainty in the mass measurements is in the determination of $i$.  If, in the future $i$ can be determined to a better accuracy, then one can combine it with an accurate determination of $q$ to measure precise binary masses.  In this paper we determine the  system parameters of  \\target\\ (primarily $q$ and $i$) by modelling the secondary star's Roche-lobe distorted absorption-line profiles with our X-ray binary model.  We also use for  the first time the technique of Roche tomography to map the surface  inhomogeneities on the secondary star in a LMXB.  \\begin{figure} \\hspace*{0mm} \\psfig{angle=-90,width=8.0cm,file=rv_curve_new.eps} \\caption{The heliocentric radial-velocity curve of the secondary star in \\target. The solid line shows a sinusoidal fit to the data. The middle a bottom panels show the residual after subtracting an  X-ray binary irradiation model and a circular orbit fit respectively. The data have been folded on the orbital ephemeris and are shown twice for clarity.} \\label{fig:rvcurve} \\end{figure}  \\begin{table} \\caption{Log of VLT+UVES observations.} \\begin{center} \\begin{tabular}{lcc}\\hline UT Date          &  UT range  & \\# spectra \\\\ \\hline 2008 June  08 & 01:05 -- 07:39 & 43 \\\\ 2008 June  09 & 02:59 -- 04:38 & 12 \\\\ 2008 June  13 & 00:16 -- 00:51 &  5 \\\\ 2008 June  23 & 00:00 -- 05:06 & 35 \\\\ \\hline \\end{tabular} \\end{center} \\label{table:log} \\end{table} ",
        "conclusions": "We have determined the  system parameters of the X-ray transient \\target\\ using only phase-resolved high-resolution UVES spectroscopy. We first determined the radial-velocity curve of the secondary star. We then modeled the shape of the phase-resolved absorption line profiles using our X-ray binary model \\textsc{XrbSpectrum} that computes the exact rotationally-broadened phase-resolved spectrum. This approach does not depend on any assumptions regarding the shape of the rotation profile, limb-darkening coefficients or laws, or the effects of contamination from an accretion disk. We determined the binary mass ratio of 0.1755$\\pm0.0025$ which is a factor of $\\sim$10 more accurate than previous estimates. We also constrain the inclination angle to 32$^\\circ$$^{+8^\\circ}_{-2^\\circ}$. Combining these values with the results of the radial velocity study gives a neutron star mass of 1.94$^{+0.37}_{-0.85}$\\Msun, consistent with previous estimates.  We present the first Roche tomography reconstruction of the secondary star in a LMXB, which reveals the presence of a large, cool polar starspot, similar to those seen in Doppler images of rapidly-rotating isolated stars. The detection of a starspot supports the idea that magnetic braking plays an important role in the evolution of LMXBs."
    },
    "1402/1402.7315_arXiv.txt": {
        "abstract": "Aside from the grand-design stellar spirals appearing in the disk of M81, a pair of stellar spiral arms situated well inside the bright bulge of M81 has been recently discovered by \\citet{Ken08}. The seemingly unrelated pairs of spirals pose a challenge to the theory of spiral density \\ron{waves}. To address this problem, we have constructed a three component model for M81, including the contributions from a stellar disk, a bulge, and a dark matter halo subject to observational constraints. Given this basic state for M81, a modal approach is applied to search for the discrete unstable spiral modes that may provide an understanding for the existence of both spiral arms. It is found that the apparently separated inner and outer spirals can be interpreted as a single trailing spiral mode. In particular, these spirals share the same pattern speed 25.5~km~s$^{-1}$~kpc$^{-1}$ with a corotation radius of 9.03~kpc. In addition to the good agreement between the calculated and the observed spiral pattern, the variation of the spiral amplitude can also be naturally reproduced. ",
        "introduction": "The magnificent spiral arms observed in disk galaxies have been proposed to be a quasi-stationary phenomenon of a wave since the mid-1960s \\citep{Lin64,Lin66}. The wave interpretation of spirals resolves the so-called winding dilemma problem that spirals would otherwise rapidly wrap up if a spiral is made of the same material in a differentially rotating disk. The stellar spiral arms are assumed to be a small perturbation superimposed on an basic state and, therefore, can be understood with linear analysis. Within the framework of a density wave, the shape of spirals is determined by a dispersion relation, which is an intrinsic property of a self-gravitating disk. Soon in the late 1960s, the working hypothesis of quasi-stationary spiral structure (QSSS) led to the study of the non-linear gas response to the underlying stellar spirals. The bright new born stars along spirals are interpreted to be a result of gas compression as molecular clouds pass through galactic shocks \\citep{Rob69}. However, it was immediately recognized that the picture of the spiral density wave was not complete \\citep{Too69}. The spiral structures are bound to disappear quickly as the wave packets propagate away from the disk. Hence, a feedback mechanism would be required to support the QSSS hypothesis.  In 1970s, both in hydrodynamics and stellar dynamics, a self-consistent theory of global spiral mode was developed to explain the excitation and maintenance of density waves. Unlike the early development of spiral density wave theory in which the pattern speed is determined empirically or inferred from observations, the modal approach associates the growth rates and the pattern speeds to the discrete eigenvalues that are subject to boundary conditions. The spiral patterns are identified as the corresponding eigen-functions of modes. In other words, spiral arms are considered to be the self-excited unstable modes with calculated pattern speeds and growth rates. The mechanisms of a wave amplification loop have been analytically explored in great detail by several authors \\citep{Mar76b,Mar77,Lin79,Ber89,Ber96,Ber00}.  \\hhwang{Given various types of spiral galaxies, the above mentioned classical density wave theory, which views the grand design spirals as intrinsic unstable modes of disks, is by no means the only possibility for the self excitation of spirals. Spirals can be interpreted as collective responses to clumps in a density distribution \\citep{Gol65,Jul66,Too91} and understood in terms of the so-called swing-amplification mechanism \\citep{Too81}. For example, the global $N$-body simulations conducted by \\citet{Sel12} show that the growth of non-axisymmetric waves can be characterized by two phases. The first, slowly growing phase is due to the swing-amplified shot noise that propagates inwards and is absorbed at the inner Lindblad resonance. The accumulating change in the distribution function at this resonance, owing to wave-particle interactions, leads to the second phase corresponding to the rapid growth of non-axisymmetric spirals. The exposure of the inner Lindblad resonance to incident wave trains is crucial for the operation of this mechanism. More recently, \\citet{Don13} demonstrate that apparent long-lived spirals can be a result of the non-linear development of gravitational wakes, which are initially induced by density clumps, such as giant molecular clouds. Once these wakes are formed, they may act as new perturbers to maintain the spiral forming activities. In this view, the grand design spirals are nothing more than connections between self-perpetuating wave segments so that the visually `long-lived' structures are understood in a statistical sense. This mechanism may potentially play a role in forming spiral structures recently observed in red disk galaxies with no gas \\citep{Mas10} and in dwarf elliptical galaxies in the Virgo cluster \\citep{Lis06, Lis09}. Here, we do not attempt to complete the list of possible mechanisms as a more thorough discussion on this topic has been recently reviewed by \\citet{Sel13}. However, we believe that different theories or mechanisms may operate in different types of spiral galaxies during secular evolution. Distinguishing between different possibilities for a specific galaxy can only rely on qualitative and quantitative comparisons between theoretical predictions and observations. In the following, we apply the modal analysis to the galaxy M81 within the framework of spiral density wave theory to confront the observations.}  M81 (NGC 3031) is one of the most well-studied spiral galaxies due to its proximity to the Milky Way and its good inclination, which is suitable for extracting kinematic information. The presence of a central bright bulge together with the nearly bi-symmetric spiral arms found in the galactic disk make M81 one of the best targets for investigating the origin of spiral structure. \\citet{Vis80b} applied the dispersion relation to a basic model for M81 to fit the shape of stellar density wave observed outside the radius of 3.5 kpc. The pattern speed, $\\Omega_p=18$ km s$^{-1}$ kpc$^{-1}$, was chosen based on the morphology of spirals and also artificially constrained by the computational consideration for the non-linear gas response discussed in the second paper \\citep{Vis80a}. In this particular model, density wave theory failed to explain structures situated inside the inner Lindblad resonance (ILR), which is located at 2.5 kpc.  \\citet{Low94} revisited the problem of spiral density waves in M81 based on the modal approach. A set of formulae was proposed to fit the observed rotation curve and the mass distribution of the disk in the range of $4<r<15$ kpc. The authors interpreted the variation of the spiral amplitude observed in the {\\it I}-band image \\citep{Elm89} as a result of the interference of trailing and leading waves that are propagating in the opposite directions. In their best-fit model, the pattern speed is determined to be $\\Omega_p=24$ km s$^{-1}$ kpc$^{-1}$ with the corresponding corotation radius located at 9 kpc. The pattern speed is in good agreement with value, $\\Omega_p=26$ km s$^{-1}$ kpc$^{-1}$, determined by the  non-linear response of gas at the 4:1 resonance \\citep{Elm89} as well as with the value, $\\Omega_p=23.4 \\pm 2.3$ km s$^{-1}$ kpc$^{-1}$, obtained by integration of the continuity equation \\citep{Tre84,Wes98}.  However, the presence of the inner turning point at 4 kpc again rules out the possibility of wave propagation in the central region $r<4$ kpc. As also cautioned in \\citet{Low94}, their analysis might be biased by the lack of reliable data in the red and infrared wavelengths at that time.  Recently, \\citet{Ken08} used {\\it Spitzer} IRAC 3.6 and 4.5 $\\mu$m near-infrared data from the Spitzer Infrared Nearby Galaxies Survey(SINGS), optical $B$, $V$ and $I$ and Two-Micron All-Sky Survey $K_s$-band data to trace the spiral density waves in M81. Using the data of different wavelengths enabled the authors to produce maps of stellar mass surface density and remove possible contaminations arising from young stars and from polycyclic aromatic hydrocarbons (PAHs). The non-axisymmetric structures (residuals) are then extracted by subtracting the axisymmetric component from the original maps. As the residual map of IRAC 3.6~$\\mu$m shown in Fig.~\\ref{fig:M81_residual_map}, aside from the outer stellar spirals, which are usually the primary focus in previous studies of M81, another pair of stellar spirals is revealed well inside the bright bulge of M81. As already noted in \\citet{Ken08}, both models as discussed in \\citet{Vis80b} and \\citet{Low94} fail to explain the presence of the inner stellar spirals. In these theoretical investigations, the bulge region is considered dynamically hot and is, therefore, hostile to the presence of density waves.  Furthermore, as also shown in \\citet{Ken08}, a good correlation is found between the inner dust spirals (IRAC 8~$\\mu$m) and the inner stellar spirals (IRAC 3.6~$\\mu$m). Since within the framework of density wave theory where density waves in the gas are allowed to propagate through the inner Lindblad resonance, \\citet{Ken08} interpreted the stellar spirals as a result induced by the gas spirals. However, from the dynamical point of view, it is unlikely that the inner stellar spirals can be induced by the gas spirals for the following reasons. First, the velocity dispersion in stars is much larger than that of gas in the bulge region. If we adopt the values from \\citet{Low94}, the radial velocity dispersion in stars can be easily more than 40 km s$^{-1}$ (cf. for the central region of the Milky Way) in the central region in contrast to the value $\\approx 10$ km s$^{-1}$ for the gas. Since the energy involved is proportional to the square of velocity dispersion, transferring energy from the gas to stars is not effective in perturbing the stellar system.  Second, little amount of gas is found in the bulge compared to the stellar mass. The total mass of the widespread HI in M81 is estimated to be $2.6\\times 10^9$ ${\\rm M}_{\\odot}$ \\citep{Yun94} in contrast to the more concentrated bulge mass enclosed within 3~kpc of $1\\times 10^{10}$ ${\\rm M}_{\\odot}$ \\citep{Low94}. Third, the thickness of the gas disk is much thinner than that of the stellar disk. In other words, the density waves in gas only effectively influence stars that are close to the mid-plane. Fourth, the absence of an ILR as shown in \\citet{Wes96} does not support the possibility of wave excitation in gas at the radius of resonance. The non-linear gas response is basically forced by the underlying stellar spirals and belongs to the type of waves discussed in \\citet{Rob69} (see also \\citet{Fuj66,Shu73,Woo75}).  The models proposed by \\citet{Low94} and \\citet{Vis80b} are not compatible to the presence of the newly discovered inner spiral structure.  With little details regarding the inner stellar spirals given in the literature, the origin and the property of the inner stellar spirals is elusive. Despite these difficulties, the faint tails of the inner spirals as shown in Fig.~\\ref{fig:M81_residual_map} hint that the inner and outer spirals may be a single spiral mode. The seemingly separate structure is simply due to low density contrast if one compares the amplitude of the spirals to the bright axi-symmetric background. If this is a plausible interpretation, the following can be inferred. First, the inner and outer arms share the same pattern speed, which has been more or less constrained in the literature for the outer spirals. Second, from a dynamical consideration, the ILR may be absent as discussed in \\citet{Adl96}, and the turning point has to be placed roughly at $r=1.4$~kpc. The location of the turning point as so chosen in \\citet{Low94} is somewhat driven by the requirement of wave behavior, i.e., the wave is not allowed to propagate in the region of $r<4$~kpc. In general, the presence of a turning point is not guaranteed for a given rotation curve. Third, the coincidence of gas(dust) and stellar spirals is a natural outcome of the theory of spiral density waves. The non-linear gas response to the underlying stellar potential is an active subject in the field.  Unlike the work of \\citet{Low94}, who use heuristic formulae to fit the rotation curve of M81 for the range $r>4$~kpc, we start with a three-component model associated with a bulge, a dark halo, and an active stellar disk. The mass model is constructed to fit the whole range of the observed rotation curve and to fit the observed mass distribution. The resultant rotation curve is then treated as input for the modal analysis. With this procedure and proper boundary conditions, we obtain unstable spiral modes that may explain the existence of spirals observed in M81. This paper is structured as follows. The three component model and the procedure of modal analysis are detailed in \\S 2. In \\S 3, we present the results obtained from the modal analysis and its comparison with observational data.  The results are discussed in \\S 4 and we summarize in the last section. ",
        "conclusions": "In this work, we interpret the stellar response of a mass model for M81 as an unstable spiral mode in a zero-thickness fluid disk. The component that actually participates in the wave propagation is the active stellar disk. The bulge and the dark matter halo components are treated as static structures to support the observed rotation curve. \\citet{Low94} introduce a mass reduction factor $f$ for an exponential disk to incorporate the effect of gravity dilution due to the disk thickness. Their active disk density profile is therefore given independently of the rotation curve. Despite the fact that the factor $f$ originates from the dynamical consideration, we have not introduced this additional adjustable function in modeling the non axisymmetric structure for two reasons. First, it is well-known that decomposing the surface light profile into different mass components is somewhat ambiguous especially for the bulge region, since the formation process of bulges (or pseudo-bulges) and the origin of exponential disks remain unclear. Second, a degeneracy exists in using parameters for three-dimensional models in fitting observed two-dimensional data. Thus, we simply adopt $f=1$ to minimize the degrees of freedom in the modelling process. In fact, the coefficients ($\\alpha_1, \\beta_1,\\gamma_1$) that lead to departures of the $Q$-profile from unity implicitly include information of the disk thickness.  Since M81 is not a gas-rich galaxy, we have neglected the limited impact of gas on the stellar dynamics in our analysis. However, the dissipative gas may play important roles in maintaining stellar density waves. In addition to facilitating the formation of new stars through shock compression \\citep{Rob69}, the interstellar medium also serves as a coolant that saturates the amplitude of the growing stellar spiral and sustains the conditions for the occurrence of grand-design spirals for a relatively long time \\citep{Rob72,Ber89,Con86,Pat94}. Furthermore, since the density waves in the gas can be transmitted beyond the outer Lindblad resonance, the conservation of wave action may lead to the spirals in atomic gas that lie well beyond the optical disk \\citep{Ber10}. The incompleteness of our model is related to the relatively low sound speed calculated for the outer 4:1 resonance, where the surface density of gas may be significant to raise the sound speed to the level required for it to be gravitationally stable. Since the equations in our study do not include the dissipative nature of gas as well as the interaction between gas and stars, we leave the sound speed as assumed.  The radius of corotation of the spiral pattern in M81 has been determined through various methods as summarized in the Section 5.2 of \\citet{Ken08}. We do not repeat them here, but comment on the method used by \\citet{Ken08}, which is based on the phase offset between the stellar spiral and 8-$\\mu$m emission. It is known from asymptotic analyses as well as from numerical simulations that there exists a angular shift between the peak of the gas shock and the potential minimum of the stellar arm \\citep{Rob69,Vis80a,Git04}. Here, we have implicitly assumed that the dust emission at 8-$\\mu$m traces the shock locations. \\citet{Git04} performed hydrodynamical simulations without self gravity to systematically study the non-linear gas response to imposed rigidly rotating stellar spirals. Inside corotation, the gaseous shocks are offset in the upstream direction with respect to the potential minimum. Although the actual shift depends on the strength of the spirals and the sound speed of the gas, the general trend is that the phase offset increases as the corotation radius is approached. Based on the simulation results, they conclude that for two-armed spirals the radius of corotation can be determined to within 25 per cent if the angular offset exceeds $\\approx \\pi/4$. However, the active star formation as observed along spirals suggests that the self gravity of gas is important to the gasdynamics around shocks and therefore may significantly affect the phase shift. The influence of self gravity on the determination of corotation radius based on the phase-offset method will be discussed in a subsequent paper.   `Halo-disk degeneracy' is an outstanding issue when decomposing an observed rotation curve into a luminous disk and a dark matter halo \\citep{Bershady10}. This is largely due to the uncertainty in the mass-to-light ratio and its radial dependence, which depends on the models of star formation history as well as of the initial mass function. Furthermore, even for luminous components, the bulge-disk decomposition introduces a further uncertainty when modelling the mass distribution of M81. The presence of inner spirals inside the bulge suggests the co-existence of disk and bulge in the central region. Despite these difficulties, Eq.~(\\ref{equation:wave_equation}) associates a proper disk response to a proper surface mass of an `active' disk, which might be considered as the lower limit of the disk mass or mass-to-light ratio. As shown in Table~\\ref{table:mass}, the total luminous mass of our model is about 88\\% of the maximum-disk mass, which is estimated to be $8.7\\times 10^{10}$~M$_{\\odot}$ \\citep{Ken87}. In order to have the correct stellar disk response, it is interesting to note that the mass decomposition of our model is quite similar to that of \\citet{Low94} and \\citet{Ken87}. An additional diagnostic of the mass model may be obtained from the study of the non-linear gas response to the observed stellar spiral. Given the observed spiral strength and the distribution of gas, the reaction of self-gravitating gas should be compared to the substructures as seen in IRAC-8$\\mu$m image. We will show in the second paper that only the proper strength of spirals, which is associated with a proper disk model, can have an appropriate gas response.  \\hhwang{The absence of inner Lindblad resonance associated with the pattern speeds as reported in this work and inferred from recent observations seems to rule out the spiral-generating mechanism proposed by \\citet{Sel12}. There is, however, a possibility that the inner spirals in M81 rotate at a pattern speed different from that of the outer spirals. In this case, the pattern speed of the inner spirals should be lower than 25.5~km~s$^{-1}$ ~kpc$^{-1}$. It is also possible that the inner and outer spirals are simply self-perpetuating spiral features as a result of the non-linear development of gravitational wakes \\citep{Don13}. As these two sets of spirals are virtually separated at the current epoch, it is not immediately clear what these possibilities imply for the future evolution of M81. In fact, for this particular case M81, the appearance of inner spirals in the nearly uniform rotating part of the disk also poses a challenge for the picture of swing amplification in the first place, since the important ingredients for this mechanism such as shear and epicyclic motions are no longer coherent in this region. Nevertheless, we emphasize that more than one mechanism may operate for spiral galaxies that are in different evolutionary states or environments.}  \\hhwang{Although this work is based on several assumptions on the model, three predictions obtained from this particular work can be checked observationally. First, the pattern speed of the inner spirals is the same as the outer spirals. Second, the velocity dispersion as a function of galactocentric distance is constrained by the shape and the strength of spirals. Third, the strength of inner spirals as shown in Fig.~\\ref{fig:amplitude} requires observational confirmation.}"
    },
    "1402/1402.2640_arXiv.txt": {
        "abstract": "A flux rope is a domain of concentration of the magnetic field $\\textbf{B}$. Insofar as $\\textbf{B}$ outside such a domain is considered as vanishingly small, a flux rope can be described as the core of a singularity of the outer vector potential $\\textbf{A}$, whose topological invariant is the magnetic flux through the rope. By 'topological' it is meant that $\\oint_C\\textbf{A}\\cdot\\mathrm d\\textbf{s}$ measures along any loop $C$ surrounding the flux rope the same constant flux $\\Phi$. The electric current intensity is another invariant of the theory, but non-topological. We show that, in this theoretical framework, the linear force-free field (LFFF) Lundquist model and the non-linear (NLFFF) Gold$-$Hoyle model of a flux rope exhibit stable solutions distributed over quantized strata of increasing energies (an infinite number of strata in the first case, only one stratum in the second case); each stratum is made of a continuous set of stable states. The lowest LFFF stratum and the unique NLFFF stratum come numerically close one to the other, and match with a reasonable accuracy the data collected by spacecrafts travelling across magnetic clouds. The other LFFF strata do not match these data at all. It is not possible at this stage to claim which model fits better the magnetic cloud data. We also analyze in some detail the merging of tubes belonging to the same stratum, with conservation of the magnetic helicity, and the transition of a tube from one stratum to another one, which does not conserve magnetic helicity.% ",
        "introduction": "}\\label{s1} It has been shown in \\cite{kleman14} that a {\\it  {singular irrotational} magnetic field} can behave as a source for twisted flux tubes (flux ropes). The singularity we have in mind carries an electric current  $\\Ical =(c/4\\pi)\\oint_C\\Bvec\\cdot\\mathrm d\\svec$, where $C$ is any loop surrounding the flux rope. A full analysis of such an electromagnetic configuration relies on the relationship between the topology of a three-dimensional domain $\\Omega$ and the calculus of divergenceless vector fields $\\Bvec$ defined in such a domain, under the condition that $\\Bvec_n =0$ on $\\partial\\Omega$ (tangential boundary conditions); these {\\it harmonic knots} (so qualified because $\\Delta \\Bvec=0$, except on the singularities) are described in mathematical terms in \\cite{cantarella02} (where more general vector fields are also described). How the theory applies to physical situations (magnetic fields in space, in the shape of flux tubes) is summarized in Sect.~1.1; the notions to be developed in this article are introduced at the end of this section, para. c) and d). Sect.~1.2 summarizes the content and main results developed in this article.\\bigskip  1.1$-$ {\\sf Main concepts in the singularity theory of flux ropes}\\smallskip  a)$-$ \\textsf{The current intensity as the topological invariant}: In this article we shall restrict our considerations to the most simple harmonic knot, namely a domain $\\Omega$ made of the whole 3D space pierced by an infinite cylindrical hole $\\Tcal$ along the $\\bf z$-axis. According to the fundamental properties of harmonic knots \\cite{kleman14,cantarella02}, the irrotational field $\\Bvec$ is uniquely determined by the following data  i$-$ boundary conditions \\begin{equation} \\label{e11} \\Bvec_n = 0 \\ \\mathrm{on}\\ \\partial \\Tcal\\qquad\\Bvec = \\{0, 0,B_o\\} \\ \\mathrm{at \\ infinity} \\end{equation}  We shall use cylindrical coordinates all along, such that $\\Bvec=\\{B_r, B_\\theta,B_z\\}$.  ii$-$ a line integral $\\beta$ along a closed curve $C$ surrounding $\\Tcal$ \\cite{kleman14,mahajan98} \\begin{equation} \\label{e1} \\beta = \\oint_{C}\\Bvec\\cdot\\mathrm{d}\\svec. \\end{equation} % By Stokes' theorem, the quantity $\\beta$ takes the same value if $C$ is replaced by a homologous loop $C'$ ($\\Bvec$ harmonic implies $\\rot\\Bvec=0$): we call $\\beta$ a {\\sf topological invariant}. In the limit where the radius of $\\Tcal$ is reduced to a line $L$ (which we choose for simplicity to be along the $\\bf z$-axis), $\\Bvec$ becomes singular on $L$, and we may write \\begin{equation} \\label{e2} \\rot\\Bvec= \\beta\\bf{\\hat{z}}\\delta^2_{\\textit{L}}(\\rvec). \\end{equation} $\\Bvec$ can be expressed as the gradient of a multivalued scalar field $\\psi$, $\\Bvec=\\grad\\psi$ \\begin{equation} \\label{e44} \\psi= \\frac{\\beta}{2\\pi}\\theta + \\gamma z % \\qquad \\Bvec= \\{% 0,     \\frac{ \\beta}{2\\pi  r}, \\gamma\\}, \\end{equation} where $\\beta$ is related to the electric current along $L$: $\\beta=({4\\pi}/{c})\\Ical$ cf. Eq.~\\eqref{e1}, $\\gamma = B_o$. The level sets $\\psi=constant$ are ruled helicoids of pitch $-b$, with\\footnote{a right helix carries a positive pitch, a left helix a negative pitch, by convention, which choice implies the use of a right-handed coordinate frame.} \\begin{equation} \\label{e3} b= \\frac{\\beta}{\\gamma}=\\frac{4\\pi}{c}\\frac{\\Ical}{B_o}. \\end{equation} Thereby one can write \\begin{equation} \\label{e4} \\Bvec= \\{% 0,     \\frac{ B_o b}{2\\pi  r}, B_o\\} \\end{equation} \\indent These properties are {reminiscent of the theory of {\\it defects} in condensed matter physics} \\cite{klfr08}; $L$ has the status of a vortex line endowed with a helical structure $-$ which character makes it also akin to a screw dislocation line \\cite{kleman14} $-$ a typical defect defined by a scalar invariant, here $\\Ical$. In the example above, $L$ is an infinite straight line, but it can be any closed loop or any curve going both ends to infinity, that carry a {topological invariant} defined by a line integral as Eq.~\\ref{e1}, taken along any loop $C$ surrounding the singularity and not enclosing another singularity \\begin{equation} \\label{e77} \\frac{4\\pi}{c}\\Ical = \\oint_C\\Bvec\\cdot\\mathrm d\\svec. \\end{equation}  b)$-$ \\textsf{The core:} The foregoing captures the essential topological properties of an harmonic knot, but it is of course an impossible physical object, due to the presence of the $1/r$ behaviour in \\eqref{e4} (equivalently, the presence of a delta function in \\eqref{e2} tells us about this impossibility). This can be cured by drilling a cylindrical hole along $L$, thus coming back to $\\Omega$; the singularity, which obeys \\eqref{e4}, is then virtual in the vacuum of the hole. Another possibility is to fill the hole with a magnetised plasma carrying an electric current $(c/4\\pi)\\beta$, $-$ we shall call this region the {\\sf core} of the singularity $-$ such that some boundary conditions (to be discussed soon) are satisfied at the contact between this {core} $\\Tcal$ and the outside medium where $\\Bvec_o$ obeys Eq~\\ref{e4} (from now on, we shall use $\\Bvec_o$ to denote the field outside the core, $\\Bvec_i$ the field inside the core).  This configuration provides a model for a \\textit{twisted tube} (a flux rope), because the boundary conditions, whatever they may be, necessarily transmit to the core the helical nature of $\\Bvec_o$. This can take on at least two aspects: either the magnetic field $\\Bvec$ is continuous across $\\partial\\Tcal$, or else the magnetic field is allowed to be discontinuous (which requires the presence of a surface current for $r = r_0$). In both cases the vector potential $\\Avec$ is continuous, cf \\cite{griffiths99} (a discontinuity of the vector potential would imply a $\\delta$-function singularity in the $\\Bvec$-field on the boundary (see the discussion in \\cite{kleman14}, footnote 2), which is most probably not physical, and will not be considered.)\\smallskip  c)$-$ {\\sf Quantized strata}: There are {\\it restrictions} on the allowed field configurations; they have indeed to minimize the magnetostatic energy $\\Ecal = \\frac{1}{8\\pi}\\iiint_\\Omega (B_\\theta^2 +B_z^2)\\mathrm dV$ $-$ under the assumption that the field 'carries' a topological invariant ($\\Ical$ in the case above, Eq.~\\eqref{e77}) $-$  and that the boundary conditions (continuity of $\\Bvec$ for $r= r_0$ in the first case above; at infinity $\\Bvec= \\{0,0,B_o\\}$) are obeyed. The calculations yield critical points satisfying  i$-$ $\\partial \\Ecal / \\partial r_0=0$ (the condition of stability $\\partial^2\\Ecal/\\partial r_0^2 >0$ has to be checked)  ii$-$ ${B_{i,\\theta}}/{B_{i,z}}(r_0) = {B_{o,\\theta}}/{B_{o,z}}(r_0)= {b}/{2\\pi r_0}$ (boundary conditions).  One gets, for a given current intensity $\\Ical$, separated sets of solutions, {\\sf quantized strata} so to speak: the tube radii, the pitches, the helicities and the energies are distributed continuously over each of these sets. See \\cite{kleman14} and what follows for examples.  The quantization rules depend on the core model, {\\it e.g.}, linear force-free field, uniformly twisted field (non-linear), constant current density, linear azimuthal current, $\\dots$ \\cite{dasso05}. Here, we investigate the first two models. \\smallskip  d)$-$ {\\sf The flux as the topological invariant, the AB model}: We have dwelt on the case where the topological invariant is the current; this requires that  the outside field be irrotational, but non-vanishing. The assumption of irrotationality is often made, but even more often is it assumed in the numerical simulations that $\\Bvec_o$ totally vanishes, which makes simulations easier and is justified by the extremely small field measured outside the ropes, at the limit of measurement possibilities \\cite{stenflo78}. Recent measurements made by spacecrafts outside flux ropes % look like fluctuations % or perhaps flux ropes of small radius and small flux: the field obviously concentrates in the ropes themselves. Therefore it is tempting to consider that the relevant variable is no longer the magnetic field $\\Bvec$, but the {\\sf vector potential} $\\Avec, \\ \\Bvec = \\rot\\Avec$, irrotational outside the core. Aharonov and Bohm \\cite{aharonov59} have recognized that this configuration yields interesting quantum effects involving interference patterns between electron currents passing next to a (small $-$ in their case) solenoid; we do not expect to observe such effects in space, whose evidence would unambiguously prove, if they exist, that a space flux rope is a physical realization of a AB (Aharonov$-$Bohm) model.  In the Coulomb gauge $\\grad\\cdot\\Avec =0$ one has $\\Avec_{o,n} =\\Avec_{i,n} =0$, because of the continuity of the vector potential at a boundary \\cite{griffiths99}. Inside the tube the magnetic field $\\Bvec_{i}$ is continuous, outside $\\Bvec_{o}=0$; the magnetic field discontinuity at the boundary requires a surface current that compensates for the volume current inside; the total current has to vanish \\cite{kleman14}. With such a picture the relevant \\textit{topological invariant} is no longer the intensity $\\Ical$ but the flux $\\Phi$. \\begin{equation} \\label{e5} \\Phi=\\oint_C\\Avec\\cdot\\mathrm d\\svec. \\end{equation}  The concept of core is of course relevant here as in the previous case.  We shall see that in the minimization process the inner total current intensity (equal $-$ and opposite in sign $-$ to the current flowing along the boundary) still plays a fundamental role. This is discussed in Sect.~\\ref{s12} in connection with the case when the outside magnetic field is small compared to the inside magnetic field and tends to zero smoothly, {\\it i.e.,} a configuration where the topological invariant $-$ the total current intensity as long as $\\Bvec_o\\neq 0$ $-$  turns suddenly to the other topological invariant, the flux, when $\\Bvec_o=0$. At the same time the field configuration changes \\textsf{continuously}. This is a further justification of the AB approach: \\textsf{the domain outside a flux rope}, little magnetized as it is, \\textsf{can be treated as a perturbation of the AB model, i.e., $\\Avec_o$ irrotational}. \\bigskip  1.2$-${\\sf Content of this article}\\smallskip  As in \\cite{kleman14} we require the Lorentz force to vanish at any point of the core; such a force-free field obeys the Beltrami equation \\begin{equation} \\label{e99} \\ell \\rot \\Bvec = \\Bvec, \\end{equation} where $\\ell$ is a constant on each line of force of the magnetic field, since $\\grad\\cdot\\Bvec=0$. The flux rope will be supposed to be a circular cylinder with translational symmetry along its axis.  We explore two situations where the topological invariant is the {\\it flux}, {\\it i.e.,} the vector potential $\\Avec_o$ is irrotational, $\\Bvec_o =0$:\\smallskip  i$-$ the Lundquist {\\it linear} force-free field case (LFFF) \\cite{lundquist50} $-$ the linearity implies that $\\ell$ is a constant over all the rope $-$. In Sect.~\\ref{s12} we show that the Bohm$-$Aharonov model is the continuous limit of a model where $\\Bvec$ is discontinuous for $r=r_0$, when $B_o\\rightarrow 0$ Sect.~\\ref{s2}. Thus the AB model is yet an excellent approximation of a true flux rope.  ii$-$ the Gold$-$Hoyle non-linear force-free field case (NLFFF) \\cite{gold60}, Sect.~\\ref{s3}, where $\\ell$ depends on radius.\\smallskip  In the same two sections the processes of tube merging and tube transitions are analyzed. Flux ropes belonging to some stratum tend to merge while keeping in the same stratum and thus release large amounts of energy; a rope belonging to a higher stratum release some energy when falling to the lowest stratum. These two processes are very different: the merging of a collection of ropes proceeds quasi continuously at constant magnetic helicity; the transition of a rope from a stratum to another involves topological modifications which do not conserve magnetic helicity.\\smallskip  In Sect.~\\ref{s4} we compare some observational data for {\\sf magnetic clouds} \\cite{lepping97,dasso03,demoulin14,dasso05} to the theoretical results developed in Sect.~\\ref{s12}, \\ref{s2} and \\ref{s3}. {Magnetic clouds (MC)} are flux ropes that extend between the Sun and the Earth, and which are triggered by the corona mass ejections (CME) that originate at the surface of the Sun. It turns out that the lowest stratum, in the LFFF case, seems to match the observational results with a good accuracy, whereas all the higher strata do not. In the NLFFF case there is only one stratum, which is very siilar to the LFFF lowest stratum. Thus it appears difficult to decide which model suits better the observations, but it is already a remarkable fact that the observations are so well fitted, and that no higher stratum than the lowest one is observed, a result that also transpires from the theory. ",
        "conclusions": "}\\label{s4}  5.1$-$ \\textsf{{An extended summary of this article}}  \\smallskip The foregoing dwells essentially upon the theoretical properties we expect for the Lundquist and Gold-Hoyle flux ropes, in the framework of a singularity model inspired in part by the theory of defects in condensed matter physics (CMP). An important concept borrowed from CMP is that of \\textsf{topological invariant}, which is in our case either the current intensity $\\Ical$ flowing through the tube, or the magnetic flux $\\Phi$ it carries; in the first case the topological invariant is related to a singularity of an outer irrotational magnetic field $\\Bvec_o$ ($\\frac{4\\pi}{c}\\Ical=\\oint_C \\Bvec_o\\cdot\\mathrm d\\svec$), in the second case to a singularity of an outer irrotational vector potential $\\Avec_o$ ($\\Phi=\\oint_C \\Avec_o\\cdot\\mathrm d\\svec$), $C$ any loop surrounding the flux tube, the {\\sf core} of the singularity in CMP parlance. The topological invariant is accompanied by another, \\textit{non-topological}, invariant: $\\Phi$ if the topological invariant is $\\Ical$, $\\Ical$ if the topological invariant is $\\Phi$. This article is about the $\\Phi$ topological invariant.  The field configurations of the singularity {cores} (the flux ropes themselves) are not different from those directly obtained without reference to their singular origin $-$ incidentally the flux ropes don't have to be singular, they are extended modes of the original singularities $-$, but due to this origin, they can be classified into {\\sf quantized strata}. Each stratum is defined by two quantum numbers, in the LFFF case the scalars $\\eta = r_0/\\ell$ and $\\zeta = b/(2\\pi r_0)$, where $r_0$ is the radius of the rope, $\\ell$ the Beltrami constant, $b$ the pitch or the period along the tube axis (the inverse of the twist); in the GH NLFFF case there is only one stratum. In each stratum the parameters defining the flux rope ($r_0, \\ell, b, B$) can vary continuously, providing $\\eta$ and $\\zeta$ are fixed. A remarkable fact is that the first stratum $\\jmath=1$ of the LFFF flux ropes (corresponding to the smallest energies) and the unique stratum of the NLFFF flux ropes are defined by quantum numbers that are numerically very close. From the analysis in Sect. 5.2, which compares the data inferred from observations to our theory, emerges a strong feeling that all those data belong either to the LFFF first stratum, or to the unique NLFFF stratum. This is what the theory would indeed suggest. Refined measurements, in particular of $\\zeta$, should mark the difference between LFFF and NLFFF.  The classification into strata plays a leading role in the merging of a collection of ropes and in the transition of a rope from one stratum to another. A well-known property used in the analysis of reconnection is that the {\\sf magnetic helicity} (under its restricted form \\cite{berger84}) is conserved.\\footnote{We give a closed expression of the restricted helicity for a cylindrical tube, which is made possible by the quantized stratum flux rope presentation we use, Eq.~\\eqref{e288}.} We  advance that this is true for merging processes taking place in the same stratum. Merging releases a large amount of energy. In a sense, this conservation law means that the merging takes place quasi continuously. On the other hand, a process that yields a transition from one stratum $\\jmath =k'$ to a lower one $\\jmath =k\"<k'$ necessitates a change of topology $-$ but at constant topological invariant $-$ if one expects some energy release; magnetic helicity is not conserved. The relation between the magnetic topology we are alluding to and magnetic helicity, although often discussed, remains according to us an open problem.\\smallskip  In the next subsection we argue that the singularity theory is justified. The last subsection compares observational data to our theoretical predictions; we believe that this comparison supports our findings.\\bigskip  5.2$-$ \\textsf{{Flux ropes are singularities}}\\smallskip  It is a rather remarkable fact that the observation of flux tubes indicate that the magnetic flux lines are twisted, hence their denomination of \\textit{flux ropes}. Because of the boundary conditions, whatever they may be, the outside field ($\\Bvec$, or $\\Avec$, or both) must also be twisted and, in the assumption of irrotationality, their level sets must be helical (as Eq.~\\ref{e3}) and thus behave in $1/r$, consequently singular for $r=0$. This is a very strong property of flux ropes, which is difficult to remove, because of their twisted character. \\bigskip  5.3$-$ \\textsf{{Comparison to observations of magnetic clouds}} \\smallskip  {\\it In situ} observations in the interplanetary space provide a number of data on flux ropes, including magnetic clouds. These data have been thoroughly analysed by a number of authors, in the framework of various flux rope models (Lunquist, Gold$-$Hoyle, non force-free models, $\\dots$); thus we have estimates of $r_0$, $\\ell$, $B_i$ and magnetic helicities in these various models, which we can compare to the present theory. Our discussion relies mostly on the data reported in ref. \\cite{dasso05,lepping97,dasso03}.  a)$-$ {\\textsf {LFFF}} We pick three cases from the references above, where the raw data are analysed in the framework of the Lundquist model:  $-$ In \\cite{lepping97} a flux rope observed by Wind on 18 Oct. 1995 is characterized by the values  $B= 24.3 \\ \\mathrm{nT}, \\ r_0 = 0.135 \\ \\mathrm{au}, \\ \\alpha = 1/\\ell =20 \\ \\mathrm{au^{-1}}$.\\footnote{au = astronomical unit (approximately the length of the semi-major axis of the Earth's elliptical orbit around the sun)} These data yield $\\eta = r_0/\\ell \\approx 2.7$.  $-$ In \\cite{dasso03}, the data collected for a hot tube observed by Wind on Oct. 24-25 1995 are as follows $B = 7.2 \\ \\mathrm{nT}, \\ r_0 = 0.035 \\ \\mathrm{au}, \\ \\alpha = 1/\\ell =65.8 \\ \\mathrm{au^{-1}}$. These data yield $\\eta \\approx 2.303$.  $-$ The 'small' magnetic clouds reported in \\cite{lepping97} yield $B$ in the range 13.8 - $15.9 \\ \\mathrm{nT}, \\ r_0 \\approx 1.6 \\times 10^{-2}\\ \\mathrm{au}, \\ \\alpha = 1/\\ell$ in the range 100 - $170 \\ \\mathrm{au^{-1}}$. Thus the value of $\\eta$ stand between 1.63  and 2.72.  {\\sf All these values of the non-dimensional parameter $\\eta$ are quite close to $\\eta_1$, table 1, and in any case closer to $\\eta_1$ than to $\\eta_2$.} % This is all the more remarkable that the dispersion in the data ($r_0,\\ \\ell$) deduced from the observations is rather large. Notice also that there seems to be no correlation between the values of the magnetic field $B_i$ and those of $r_0$.  Thus, if one adopts the value $\\eta=\\eta_1=2.1288$ as a best fit, then $\\zeta= 3.7610$, $b=2\\pi r_0$ $=0.827$ au for {\\it e.g.} $r_0 = 0.035 \\ \\mathrm{au}$, etc. Let us keep to this latter case,  $r_0 = 0.035 \\ \\mathrm{au}$; other quantities can be calculated, as the \\textit{energy per unit length} $\\Ecal|_{r_0}$ \\eqref{e38b}, which amounts to $\\Ecal|_{r_0}=0.518 \\times10^4$ J/m, or the \\textit{current intensity} $\\Ical=\\frac{c}{4\\pi}B b= 2.1323\\times 10^{10}$ SI units, {\\it i.e.} a current density $j_z=\\frac{\\Ical}{\\pi r_0^2}=0.2462$ nT/s ($\\pi r_0^2 = 8.66\\times 10^{19}$ m$^2$). % \\smallskip  b)$-$ {\\textsf {NLFFF}} Here the raw data are analysed in the framework of the Gold$-$Hoyle model \\cite{dasso03} for the Wind Oct. 24-25 1995 event.  One gets $r_0 = 0.035$ au, $2\\pi \\varpi =46.2$ au$^{-1}$, thus $q= 2.6147$, which is remarkably close to our theoretical value $q= 2.5129$, eq.~\\eqref{32c}.  This discussion can be extended to the value of the current density $j_z$. With $q=2.5129$ (the theoretical value) as above and taking for the other quantities the values from the data analysis, one gets, employing the relation $(\\pi r_0^2) j_z = \\frac{c}{4\\pi\\varpi}B_z(0) \\frac{q}{1+q}$, $j_z = 0.0302$ nT/s, which is 8 times smaller than the value obtained for the same MC in the LFFF model \\smallskip  {\\textsf {Remark}} It is worth comparing these calculated values to the value of $j_z$ in the {\\it third} model analysed in \\cite{dasso03}, (a non force-free field model with constant current density) the only model that allows for an estimation of $j_z$ directly from the raw data. It appears that the favoured value is $j_z= 0.12$ nT/s, which is very different from what our NLFFF calculation above; on the other hand this estimation of $j_z$ is closer to the prediction made from the LFFF model."
    },
    "1402/1402.0968_arXiv.txt": {
        "abstract": "CMB spectral distortions are induced by Compton collisions with electrons. We review the various  schemes to characterize the anisotropic CMB with a non-Planckian spectrum.   We advocate using logarithmically averaged temperature moments as the preferred language to describe these spectral distortions, both for theoretical modeling and observations.  Numerical modeling is simpler, the moments are frame-independent, and in terms of scattering the mode truncation is exact. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1475_arXiv.txt": {
        "abstract": "The long residence times and small anisotropies of cosmic rays suggest that they are well confined and well scattered by the Galactic magnetic field. Due to the disklike shape of the confinement volume, transport in the vertical direction, perpendicular to the mean Galactic magnetic field, is key to cosmic ray escape. It has long been recognized that this vertical transport depends both on the vertical component of the fieldlines themselves and on the extent to which the cosmic rays are tied to the fieldlines. In this paper we use magnetic fields with very simple spatial and temporal structure to isolate some important features of cross fieldline transport. We show that even simple magnetic nonuniformities combined with pitch angle scattering can enhance cross fieldline transport by several orders of magnitude, while pitch angle scattering is unnecessary for enhanced transport if the field is chaotic. Nevertheless, perpendicular transport is much less than parallel transport in all the cases we study. We apply the results to confinement of cosmic rays in the Fermi Bubbles. ",
        "introduction": "\\label{sec:introduction}  It is well accepted that cosmic rays below $\\sim 10^{18}$ eV are confined to the Milky Way by the Galactic magnetic field. The confinement volume is a thick disk, suggesting that the primary direction of escape is perpendicular to the Galactic plane. Furthermore, the distribution of cosmic ray arrival directions at the Earth is nearly isotropic. These features, together with observations of the Galactic magnetic field which show it to be oriented preferentially parallel to the Galactic plane but with a substantial randomly oriented component, have led to a picture in which cosmic rays are isotropized by scattering from small amplitude, short wavelength magnetic fluctuations and achieve their vertical displacement both by following the magnetic fieldlines as they snake their way out of the Galactic plane and by transport across the field itself (see \\cite{hillas_2006, aharonian_2012, zweibel_2013} for some recent reviews).  This spatial transport with respect to the background medium (as opposed to, e.g., advective transport by a Galactic wind) is frequently described by a diffusion tensor $\\bkappa\\equiv\\kappa_{\\perp}\\bI - (\\kappa_{\\perp}-\\kappa_{\\parallel})\\bB\\bB/B^2$, where $\\bI$ is the unit tensor and $\\bB$ is the {\\textit{mean}} magnetic field. Simple parameterizations of the diffusion tensor have led to many successful descriptions of cosmic ray properties (e.g. \\cite{strong_1998, strong_2010, vladimirov_2011}).  Estimating $\\kappa_{\\perp}$ and $\\kappa_{\\parallel}$, % however, and even establishing the circumstances under which transport is diffusive ($\\langle\\Delta x^2\\rangle\\propto\\Delta t$), as opposed to subdiffusive ($\\langle\\Delta x^2\\rangle\\propto\\Delta t^{\\alpha}$; $\\alpha < 1$), or superdiffusive ($\\langle\\Delta x^2\\rangle\\propto\\Delta t^{\\alpha}$; $\\alpha > 1$), has proven to be difficult. Transport results from the interplay of many effects: the large scale magnetic field geometry, the properties of the turbulence, the rate of scattering by small scale fluctuations, and the extent to which magnetic fieldlines diverge on scales less than the particle gyroradius. It is generally agreed that if the particles are tied to the fieldlines, the maximum perpendicular diffusion rate is set by the rate of perpendicular fieldline wandering, scaled by particle velocity (Fieldline Random Walk, or FLRW), and that parallel scattering reduces $\\kappa_{\\perp}$ below its FLRW value \\citep{jokipii_1966, jokipii_1969, rechester_1978, giacalone_1999, minnie_2009}. For a recent review of the subject, see \\cite{shalchi_2009}.  In this paper, we address a different problem: the rate at which cosmic rays cross the {\\textit{exact}} field, rather than the {\\textit{average}} field. This problem has received relatively less attention (see \\cite{barge_1984} however). While for cosmic ray confinement in the Galaxy it is primarily vertical transport (i.e. perpendicular to the mean magnetic field) which is important,  transport across the exact field matters for other problems. Confinement of cosmic rays to bubbles blown by AGN, which presumably are bounded by magnetic surfaces is one such problem. The Fermi Bubbles \\citep{finkbeiner_2004, dobler_2010} - which may have been created by sporadic activity in the Galactic Center - are a local example. Numerical simulations of bubble evolution with perpendicular cosmic ray diffusion suppressed yield surface brightness profiles which agree quite well with observations, while surface brightness profiles based on simulations with $\\kappa_{\\perp}=\\kappa_{\\parallel}$ are more diffuse than observed \\citep{yang_2012}.  Cross fieldline transport is easily estimated in two extreme cases: drifting due to large scale structure, and scattering by small scale structure. Drifts occur either when the direction or strength of the magnetic field varies in space or when electrical or gravitational acceleration $\\ba$ acts on the particles. The drift velocity $v_D$ produced by magnetic field inhomogeneity on scale $L$ can be estimated as $v_D/v \\sim r_g/L$, where $v$ is the particle speed and $r_g$ is the particle Larmor radius. The drifts induced by nonmagnetic forces are of order $a_{\\perp}/\\om_{g}$, where $a_{\\perp}$ is the component of nonmagnetic acceleration normal to $\\bB$ and $\\om_{g}$ is the relativistic gyrofrequency. If one adopts values of $L$ and $a_{\\perp}$ characteristic of global properties of the Galaxy, then the drifts are too small to have anything to do with cosmic ray escape, except for very high energy cosmic rays. However, there is a certain amount of random drifting due to moderate scale turbulence.  Scattering by small ($r_g$ scale) fluctuations is generally thought to be the mechanism responsible for parallel diffusion, but also produces perpendicular diffusion. Every pitch angle scattering event $\\delta\\theta\\sim \\delta B/B$, where $\\delta B/B$ is the fluctuation relative amplitude, is accompanied by a cross fieldline displacement of order $\\delta r\\sim r_g\\delta\\theta$. The corresponding perpendicular diffusion coefficient $\\kappa_{\\perp}\\sim v r_g$, resulting in a vertical escape time of order $L^2/(v r_g)$, is similar to what is estimated from drifts and, again, too small to have anything to do with the escape of GeV cosmic rays, where the bulk of the cosmic ray distribution lies.  If magnetic fieldlines separated by less than a fiducial gyroradius wander away from each other, particles do not follow fieldlines. This case is difficult to treat analytically, but has been studied numerically \\citep{jokipii_1973, qin_2002, hauff_2010, lazarianyan_2013}, together with its effects on diffusion across the mean field.   A realistic numerical study of cosmic ray propagation based on integration of particle orbits would include both the large scale geometry of the Galactic magnetic field and the small scale fluctuations which provide pitch angle scattering and possibly second order Fermi acceleration. Such a magnetic field model would have to span at least 10 orders of magnitude in lengthscale, while current numerical simulations of interstellar magnetized turbulence rarely capture more than two orders of magnitude in scale. And even if one prescribed a magnetic field model with a realistic range of scales, accurately integrating large enough numbers of particle orbits over long times and distances required for a statistically representative sample would be prohibitively expensive.  In this study of particle transport, we address the range of scales problem by integrating particle orbits in magnetic field structures much larger than a larmor radius with pitch angle scattering by gyroradius scale fluctuations modeled as a random process. A similar approach was followed in \\cite{neuer_2006a, neuer_2006b}, but their random process, which is intended to model collisions, does not conserve energy, while ours, which is intended to model pitch angle scattering by nearly static fluctuations, does. This dual approach allows us to probe the effects of fieldline geometry and variation of the field on large scales while retaining pitch angle diffusion, and at the same time avoiding the need to integrate over a large dynamic range in scales of magnetic structure. We find that in many cases diffusive behavior is recovered, but that the diffusivities do not depend on fluctuation amplitudes according to the prescription derived from quasilinear theory.  In \\S\\ref{sec:setup} we introduce the basic magnetic field models we use, discuss the implementation of pitch angle scattering, show how we measure diffusion, and establish upper and lower bounds on the cross field diffusivity. In \\S\\ref{sec:results} we present the results for the different field models and discuss their physical basis. Section \\ref{sec:discussion} is a discussion and comparison with other results and \\S\\ref{sec:conclusions} is a summary. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have investigated the transport of cosmic rays across magnetic fields. In particular, we considered the interplay between magnetic geometry and cross fieldline transport by pitch angle scattering. For simplicity, we chose simple, single scale 2D cellular magnetic geometries superimposed on a guide magnetic field. One case is static, with integrable fieldlines, and the other is periodic in time, with chaotic fieldlines. The scattering is assumed to be due to gyroscale magnetic turbulence, while the cell size is well above the gyroradius. Our study complements the large body of work on transport across {\\textit{mean}} magnetic fields.  We showed that the running spatial diffusion coefficients $D(t)$ defined in eqn. (\\ref{runningD}) approach steady values $\\kappa$ over time, indicating diffusive (as opposed to subdiffusive or superdiffusive) behavior. We argued that due to the bounded horizontal excursions of the fieldlines chosen for our study, $\\kappa_{\\perp}$ represents true cross fieldline transport. Cellular structure enhances transport by several orders of magnitude as compared to transport in a uniform field. We found that fields which vary on a timescale similar to the timescale on which particles orbit the cells have fast transport even without pitch angle scattering. But in all the cases we considered, perpendicular diffusion is much slower than parallel diffusion. We derived upper and lower bounds on the diffusivity and compared our results to those obtained for advection-diffusion of passive scalars, and in fully turbulent magnetic fields, but did not find the same scalings. We did, however, apply our results to particle confinement in magnetically bounded structures such as the Fermi Bubbles, and argued that escape by cross fieldline transport would be slow.  In reality, cosmic rays encounter magnetic turbulence on a vast range of scales. Each particle will be affected more by some scales than others, but the interplay of multiple scales is sure to give rise to interesting effects not encountered here. We will take up some of these effects in future work."
    },
    "1402/1402.3187.txt": {
        "abstract": "{ We present a comparison of the molecular gas properties in the outflow vs. in the ambient medium of the local prototype radio-loud and ultraluminous-infrared galaxy 4C12.50 (IRAS13451$+$1232), using new data from the IRAM Plateau de Bure interferometer and 30\\,m telescope, and the \\her\\ space telescope.  Previous \\htwo\\ (0-0) S(1) and S(2) observations with the \\spi\\ space telescope had indicated that the warm ($\\sim$400\\,K) molecular gas in 4C12.50 is made up of a 1.4($\\pm$0.2)$\\times$10$^8$\\msun\\  ambient reservoir and a 5.2($\\pm$1.7)$\\times$10$^7$\\msun\\ outflow. The new \\cooz\\ data cube indicates that the corresponding cold (25\\,K) \\htwo\\ gas mass is 1.0($\\pm$0.1)$\\times$10$^{10}$\\msun\\ for the ambient medium and $<$1.3$\\times$10$^8$\\msun\\ for the outflow, when using a CO-intensity-to-H$_2$-mass conversion factor $\\alpha$ of 0.8\\msun /(K\\kms\\,pc$^2$).  The combined mass outflow rate is high, 230-800\\msun /yr, but the amount of gas that could escape the galaxy is low. A potential inflow of gas from a 3.3($\\pm$0.3)$\\times$10$^8$\\msun\\ tidal tail could moderate any mass loss. The mass ratio of warm-to-cold molecular gas is $\\gtrsim$30 times higher in the outflow than in the ambient medium, indicating that a non-negligible fraction of the accelerated gas is heated to temperatures at which star formation is inefficient. This conclusion is robust against the use of different $\\alpha$ factor values, and/or different warm gas tracers (\\htwo\\ vs. \\htwo\\ plus CO): with the CO-probed gas mass being at least 40 times lower at 400\\,K than at 25\\,K, the total warm-to-cold mass ratio is always lower in the ambient gas than in the entrained gas. Heating of the molecular gas could facilitate the detection of new outflows in distant galaxies by enhancing their emission in intermediate rotational number CO lines. } ",
        "introduction": "\\label{sec:intro}  Mass accretion events onto black holes (BHs) can release more energy than the binding energy of gas in galaxies \\citep{fabian12}, and part of this energy can be deposited onto the interstellar medium (ISM) through radiation, accretion-disk winds, and radio jets. The energy carried by these so-called feedback mechanisms can excite or kinematically distort gas on large scales \\citep[e.g., over areas of square kiloparsecs;][]{lipari04,fu09,rupke11,westmoquette12}. It can even affect galaxies external to the active galactic nucleus (AGN) that generated it \\citep[e.g.,][]{croft06,cantalupo12}. The effects of feedback on the evolution of galaxies are most severe if part of the gas is accelerated beyond escape velocity. The high-velocity gas will then be lost to the intergalactic medium. Any star formation associated with it will be quenched. For the remaining gas, star formation will be delayed. The low-velocity gas will cool radiatively, resettle in a disk, and restart its collapse towards the formation of new cores. For these reasons, AGN feedback has been considered as a source of negative feedback, capable of regulating galaxy growth. Positive feedback, i.e., compression of gas leading to an enhancement of star formation has also been occasionally suggested \\citep[e.g.,][]{vanbreugel85,gaibler12,silk13}.  AGN feedback,  implemented either as a source of heating coupled to the hot gas or as a source of momentum coupled to the cold gas, was included in cosmological simulations to suppress star formation at earlier epochs and to make the predicted number counts of massive present-day galaxies match the observed ones \\citep{granato04,croton06,bower06,somerville08,booth09,debuhr12}.  The claim that feedback could change the high-mass end of the local galaxy mass function by more than an order of magnitude was however contested by \\citet{birnboim07}, \\citet{dekel08}, and \\citet{bernardi13}, who argued that the numbers could be revised downwards through, e.g., different stellar population modeling assumptions.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 1%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure}[t!] \\begin{center} \\includegraphics[width=\\columnwidth]{f1.eps} \\caption{ {\\it Left:} 102 GHz continuum image of 4C12.50 (in units of Jy/beam). {\\it Right:} Continuum-free \\cooz\\ line intensity map of 4C12.50 averaged over the -500 \\kms\\ to 500 \\kms\\ range, showing extended emission that is offset from the continuum peak (marked with a cross). The contours are at the 3$\\sigma$ levels for this image, i.e., at steps of 0.36\\,mJy/beam.} \\label{fig:CO_cont_line} \\end{center} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 1%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   The change in the galaxy mass function by AGN feedback remains to be quantified or constrained by observations. To quantify it, we need to determine the frequency of AGN-driven outflows that can considerably delay or quench star formation in their host galaxies. Even though high-velocity outflows have long been identified in optical and radio data of ionized and neutral atomic gas tracers in local AGN \\citep[e.g.,][]{heckman00,rupke05,martin05,morganti05,holt08}, this research field became particularly appealing thanks to the launch of infrared (IR) space missions (e.g., \\her ) and the improvement on the sensitivity and the spectral range of ground-based mm facilities, which enabled us to probe the molecular gas content of outflows.  Being the densest ISM component, the molecular gas can considerably contribute to the mass of the entrained gas, if not dominate it.  Motions of molecular gas that could be associated with a nuclear outflow have been presented for $\\sim$50 local galaxies to date \\citep{sakamoto06,sakamoto09,leon07,feruglio10,fischer10,alatalo11,rangwala11, sturm11,krips11,dasyra_combes11,dasyra_combes12, aalto12,tsai12,morganti13a,combes13,spoon13,veilleux13,cicone14}. Mass measurements of the accelerated gas were presented in $\\sim$one fifth of the cases and found to be in the range 10$^6$-10$^9$\\msun .  To identify spectral lines with strong outflow signatures, and use them for the systematic and efficient detection of new outflows, we need to perform excitation studies of the gas in the flow, motivated by the argument that the AGN feedback could leave both kinematic and excitation signatures on the ISM. Such studies have nonetheless only been presented for a couple of sources.  In Mrk 231, \\citet{cicone12} found no evidence for a strong contribution of shocks on the excitation of outflowing CO molecules at states of low rotational number $J$: the \\coto /\\cooz\\  flux ratio was comparable in the outflow and in the ambient medium.  In this paper, we present evidence of heating of the molecular gas in the outflow of the local radio-loud and ultraluminous-infrared system 4C12.50 (IRAS 13451$+$1232), i.e., a late-stage merger of two galaxies where the nuclei are separated by 4.4\\,kpc and surrounded by smaller structures \\citep{batcheldor07}.  A relativistic radio jet with an intrinsic bulk flow speed of at least 0.8$c$ emerges from the system's west nucleus \\citep{lister03}. The setup of a radio jet propagating through a gas-rich  galaxy, which is rare in the local Universe but common at high redshift \\citep[e.g.,][]{sajina08,ivison12}, makes 4C12.50 an excellent target for studying the properties of the outflowing molecular gas. The uniqueness of 4C12.50 lies in the discovery and the characterization of a warm molecular gas outflow in it: it was the only source in the \\spi\\ archive for which an outflow-related wing was detected in two purely rotational \\htwo\\ lines, providing the entrained gas excitation temperature and mass \\citep{dasyra_combes11}. Cold gas was also discovered in the outflow thanks to absorption seen in CO(2-3) with a significant optical deph against the millimeter continuum source \\citep{dasyra_combes12}. In this paper, we compare the properties of the gas in the outflow vs. that in the ambient medium using new submillimeter and millimeter data, and we discuss the implications of our findings for future searches of massive molecular gas outflows. Throughout our work, we adopt a $\\Lambda$CDM cosmology with H$_0$=70 \\kms\\  Mpc$^{-1}$, $\\Omega_{M}$=0.3, and $\\Omega_{\\Lambda}$=0.7.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 2%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\begin{center} \\includegraphics[width=8cm,height=6.7cm]{f2.eps} \\caption{Position-velocity diagram of the continuum-free \\cooz\\ emission along the east-west axis, displayed for a velocity resolution of 6\\kms . The x-axis offset is computed from the radio core location, which is marked with a dashed line. The offset increases (/decreases) towards the west (/east).} \\label{fig:pos_vel} \\end{center} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 2%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 3%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} \\begin{center} \\includegraphics[width=19.3cm]{f3.eps} \\caption{Continuum-free \\cooz\\ line intensity map of 4C12.50, displayed for $\\sim$114 \\kms\\ bins (in the range $-$315\\kms\\ to 369\\kms ). The number at the upper-left corner of each frame corresponds to the bin mean velocity. Contours are plotted at integer multiples of the 5$\\sigma$ level for these images, i.e., at steps of 2.4 mJy/beam. A shift towards the west of the line peak position from negative to positive velocities, as well as a distinct kinematic component near rest-frame velocity are observed.} \\label{fig:COmap_Vshift} \\end{center} \\end{figure*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 3%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 4%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} \\begin{center} \\includegraphics[width=19.3cm]{f4.eps} \\caption{ Same as in Fig.~\\ref{fig:COmap_Vshift}, but for bins of 17\\kms\\ (in the range 0 to 34\\kms ). Contours are plotted at integer multiples of the 3.5$\\sigma$ level for these images, i.e., at steps of 3.3 mJy/beam.} \\label{fig:COmap_structure} \\end{center} \\end{figure*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIGURE 4%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %___________________________________________________ SECTION DATA ___________________________________________________________ ",
        "conclusions": "\\label{sec:conclusions}  We used IRAM PdB, 30\\,m telescope, and \\her\\ Space Telescope data to study the  content, excitation, and kinematic properties of the molecular gas in the outflow and in the ambient medium of a local prototype radio-loud, ultraluminous-infrared galaxy, 4C12.50  (also known as IRAS13451$+$1232). We specifically looked for gas heating and expulsion signatures, i.e., characteristic of the two feedback-related processes that can delay star formation. Combining new and previous observations, we found evidence for both processes taking place, and we drew the following main conclusions.  \\begin{itemize} \\item The \\cooz\\ emission spatially consists of a newly discovered tidal tail, which extends out to 12\\,kpc south-west of the nucleus, and a main component, which is marginally resolved within the beam radius (4.2 kpc). A shift in the location of the maximum emission with velocity indicates that the main component could be tracing either two merging nuclei or a rotating disk. Its reservoir contains 1.0($\\pm$0.1)$\\times$10$^{10}$\\msun\\ of cold \\htwo\\ gas for $\\alpha$=0.8\\msun /(K\\kms\\,pc$^2$). \\item The CO SLED, which is primarily tracing the ambient gas in 4C12.50, is very similar with that of the Milky Way gas for low-$J$ numbers. Fitting of the CO SLED with RADEX indicates that the gas mass probed by CO molecules at 400\\,K is less than 0.025 times the gas mass probed by CO molecules at 25\\,K. Likewise, the \\htwo\\ gas mass probed by CO molecules at 400\\,K is less than 1.7 times the mass of the 400\\,K gas probed by the purely rotational \\htwo\\ lines in the mid-infrared. \\item In the outflow, the mass of the cold (25\\,K) \\htwo\\ gas is $<$1.3($\\pm$0.1)$\\times$10$^8$\\msun , i.e., at most twice as high as the mass of the warm (400\\,K) \\htwo\\ gas emitting in the mid-infrared. By combining the \\htwo -based mass flow rate measurement from the mid-infrared data, 230\\msun /yr, with the CO-based mass flow rate upper limit from the mm data, 570\\msun /yr, we conclude that the feedback mechanism(s) could be accelerating up to 800\\msun /yr of molecular gas for an outflow radius of 150\\,pc. \\item A low fraction of the outflowing gas ($<$30\\% when ignoring dark matter) could escape the system and be lost to the intergalactic medium. Any mass loss could be moderated or counteracted by a potential inflow of gas from the tidal tail containing 3.3($\\pm$0.3)$\\times$10$^8$\\msun\\ of cold \\htwo . The star formation in the host of 4C12.50 is likely to be delayed, but not entirely suppressed by the feedback. \\item The mass ratio of warm-to-cold \\htwo\\ gas is elevated in the outflow with respect to the ambient medium by a factor of $\\gtrsim$30 for $\\alpha$=0.8\\msun /(K\\kms\\,pc$^2$). The conclusion that the accelerated molecular gas is heated is robust against two main sources of uncertainty: the chosen $\\alpha$ conversion factor, and the mass of the warm gas probed by the CO. While the discrepancy will be lower for higher values of either parameter, there is no combination of acceptable parameter values that could make the warm-to-cold gas mass ratio in the ambient medium reach that in the outflow. \\item A scenario in which outflows could be probed by a different CO SLED than the rest of the ISM is likely. It remains to be tested with sub-millimeter observations. \\end{itemize}     %___________________________________________________ SECTION BIB ___________________________________________________________"
    },
    "1402/1402.3700_arXiv.txt": {
        "abstract": "We present evidence for the rapid variability of the high velocity iron K-shell absorption in the nearby ($z=0.184$) quasar PDS\\,456. From a recent long \\suzaku observation in 2013 ($\\sim1$\\,Ms effective duration) we find that the the equivalent width of iron K absorption increases by a factor of $\\sim5$ during the observation, increasing from $<105$\\,eV within the first 100\\,ks of the observation, towards a maximum depth of $\\sim500$\\,eV near the end. The implied outflow velocity of $\\sim0.25$\\,c is consistent with that claimed from earlier (2007, 2011) \\suzaku observations. The absorption varies on time-scales as short as $\\sim1$\\,week. We show that this variability can be equally well attributed to either (i) an increase in column density, plausibly associated with a clumpy time-variable outflow, or (ii) the decreasing ionization of a smooth homogeneous outflow which is in photo-ionization equilibrium with the local photon field. The variability allows a direct measure of absorber location, which is constrained to within $r=200-3500$\\,\\rg of the black hole. Even in the most conservative case the kinetic power of the outflow is $\\gtrsim6\\%$ of the Eddington luminosity, with a mass outflow rate in excess of $\\sim40\\%$ of the Eddington accretion rate. The wind momentum rate is directly equivalent to the Eddington momentum rate which suggests that the flow may have been accelerated by continuum-scattering during an episode of Eddington-limited accretion. ",
        "introduction": "Outflows (or winds) are likely to be a natural and unavoidable result of the accretion process (e.g., \\citealt{king03,ohsuga09}). In Active Galactic Nuclei (AGN) the feedback associated with outflowing matter is believed to play an important role in shaping the co-evolution of the central massive black hole and the host galaxy (\\citealt{king03,dimatteo05}), plausibly leading to the observed AGN--host-galaxy relationships, e.g., $M-\\sigma$ (\\citealt{king10,zubovas12,mcquillin13}). Recently, a number of massive ($\\nh\\sim10^{23}$\\,\\pcmsq), high-velocity ($\\vout\\gtrsim0.1$\\,c) outflows, as revealed through the presence of blue-shifted Fe\\,K absorption at $E>7$\\,keV in the rest-frame, have been found in luminous AGN (e.g., \\citealt{pounds03, chartas03, reeves09, tombesi10, gofford13}). The large wind velocity---which indicates an origin directly associated with the accretion disk, hence leading them to be dubbed `disk-winds'---implies that the wind may be energetically significant in terms of feedback (i.e., $L/\\lbol\\sim0.5-5\\%$, \\citealt{hopkins10,dimatteo05}).  Nonetheless, one key determination currently lacking is a direct measurement of the wind location with respect to the central black hole, which is of fundamental importance when it comes to determining the wind energetics. Previous studies of the disk-wind phenomenon have employed simple geometric and kinematic relations to constrain the location of the absorbing gas (e.g., see \\citealt{tombesi12,tombesi13} for an outline of these arguments), but these arguments can lead to considerable uncertainties on the gas location which makes it difficult to confidently determine the wind energetics. The best way to overcome these limitations it to directly determine the location of the absorbing gas by establishing how it varies over short time-scales. However, whilst such line-variability has sometimes been observed in soft X-ray grating spectra below $E=2-3$\\,keV, hence leading to robust constraints on the location of the soft X-ray warm absorber in a number of AGN, e.g., in MRK\\,509 (\\citealt{kaastra12}) and MR\\,2251-178 (\\citealt{reeves13a}), to date it has proven much more difficult to measure the necessary line variability at harder X-ray energies (i.e., $E>7$\\,keV) where modern detectors tend to be less sensitive (e.g., \\citealt{giustini11}). In this work we overcome these limitations and report the first direct constraints on the location of a high-velocity Fe\\,K disk-wind, as measured in the powerful quasar PDS\\,456.  PDS\\,456 ($z=0.184$) is a luminous ($\\lbol=10^{47}$\\,\\ergps; \\citealt{simpson99,reeves00}) radio-quiet quasar which harbours one of the most powerful Fe\\,K disk-winds currently known (\\citealt{reeves09}, hereafter `R09'). \\xmm first found the X-ray spectrum to be absorbed at $E>7$\\,keV in 2001, with the absorption attributable to highly-ionised iron (\\citealt{reeves03}). A subsequent \\suzaku observation in 2007 revealed two highly significant absorption lines at $9.08$ and $9.66$\\,keV in the quasar rest-frame (R09). These lines are most likely associated with resonant absorption from \\fexxv~\\hea and \\fexxvi~\\lya, which hence implies an outflow velocity in the range $\\vout\\sim0.25-0.3$\\,c (R09). In a more recent \\suzaku observation (2011) we again found a broad absorption trough at $\\sim9$\\,keV in the source rest-frame (Reeves et al. 2013, hereafter `R13'), and we found that the absorption in both 2007 and 2011 could be due to the same flow of gas which is in photo-ionization equilibrium with the emergent X-ray emission. However, due to the large time difference between the observations, direct constraints could not be placed on the radial location of the absorbing gas. Even so, we were able to describe the absorption in both observations using the self-consistent disk-wind models of \\citet{sim08,sim10a}, hence showing the variable wind profile to be consistent with a radiatively-driven flow launched from the inner accretion disk (R13).  In this paper we report the first results from our extensive observational campaign of PDS\\,456 with the \\suzaku, \\xmm and \\nustar satellites in 2013 (Feb--Sept). Here, we focus specifically on characterising the remarkable spectral variability exhibited by the Fe\\,K wind during the new long \\suzaku observation (Feb--Mar 2013), and use this variability to place the first direct constraints on the wind location. A full spectral analysis of both these \\suzaku data and the accompanying contemporaneous \\xmm and \\nustar campaign will be presented in forthcoming work. ",
        "conclusions": "This paper presents the first results from an extensive \\suzaku, \\xmm and \\nustar observing campaign of powerful quasar PDS\\,456, occurring between Feb--Sept 2013. We have reported on the remarkably variable high-velocity Fe\\,K-shell wind which is evident in the new long ($\\sim1$\\,Ms duration) \\suzaku observation. Consistent with earlier (2007, 2011) \\suzaku observations, the wind is again detected through absorption at $\\sim9$\\,keV in the source rest-frame ($\\vout\\sim0.25$\\,c). The absorption line depth increase by a factor of $\\sim5$ during the new observation. This variability is equally well modelled by (i) an outflow in photo-ionization equilibrium which recombines in response to decreasing source flux, or (ii) an inhomogeneous clump of gas which transits the line-of-sight to the quasar. The variability allows us to directly determine the radius to the gas, e.g., as part of a clumpy outflow, which is constrained to lie between $r\\sim10^{2}-10^{3}$\\,\\rg of the black hole. Even in the most conservative case, the kinetic power of the flow is a significant fraction ($\\gtrsim6\\%$) of the Eddington luminosity, and is comparable to the $\\sim0.5-5\\%$ of $\\lbol$ thought necessary for significant feedback. The momentum rate of the flow is equivalent to the Eddington momentum rate which is consistent with the flow being radiatively-accelerated by electron scattering during a near-Eddington-limited accretion episode."
    },
    "1402/1402.0766_arXiv.txt": {
        "abstract": "A significant proportion ($\\sim30\\%$) of the short-duration gamma-ray bursts (SGRBs) localised by {\\em Swift} have no detected host galaxy coincident with the burst location to deep limits, and also no high-likelihood association with proximate galaxies on the sky. These SGRBs may represent a population at moderately high redshifts ($z\\gtrsim1$), for which the hosts are faint, or a population where the progenitor has been kicked far from its host or is sited in an outlying globular cluster. We consider the afterglow and host observations of three ``hostless'' bursts (GRBs 090305A, 091109B and 111020A), coupled with a new observational diagnostic to aid the association of SGRBs with putative host galaxies to investigate this issue. Considering the well localised SGRB sample, 7/25 SGRBs can be classified as ``hostless'' by our diagnostic. Statistically, however, the proximity of these seven SGRBs to nearby galaxies is higher than is seen for random positions on the sky. This suggests that the majority of ``hostless'' SGRBs have likely been kicked from proximate galaxies at moderate redshift. Though this result still suggests only a small proportion of SGRBs will be within the AdLIGO horizon for NS-NS or NS-BH inspiral detection ($z\\sim0.1$), in the particular case of GRB\\,111020A a plausible host candidate is at $z=0.02$. ",
        "introduction": "Short-duration gamma-ray bursts (SGRBs)\\footnote{Usually taken as those with $T_{90}<2$\\,s, where $T_{90}$ is the time over which 90\\% of the prompt gamma-ray emission is observed.  As a class, the SGRBs also have harder spectra than LGRBs, but there is significant scatter.} reside in different environments from their long GRB counterparts (LGRBs). While, LGRBs are usually found in the brightest regions of star forming host galaxies \\citep{GRBoffset,LGRBs_in_bright_regions}, SGRBs often appear in galaxies with established older populations with a handful linked strongly to elliptical galaxy hosts \\citep{Gehrels_050509B_host,Bloom_2006_050509B,Berger_2005_050724_host}. This diversity in environments almost certainly reflects differing progenitors for the LGRBs and SGRBs \\citep{Bloom_2006_diverse_progenitors}. Indeed, in the majority of LGRBs at low redshift it has been possible to isolate the signature of a type Ic supernova, suggesting that their progenitors are Wolf-Rayet stars \\citep[][and references therin]{SN_connection}. In contrast, deep searches in SGRBs at similar redshift fail to locate any such signatures, and offer further evidence that the progenitors of short bursts are not related to stellar core collapse \\citep{Bloom_2006_050509B,Hjorth_2005_050709_OA,Rowlinson_2010_080905A_Opt}. The varied host demographics, offer part of the picture, but the locations of the bursts on their hosts are also greatly diagnostic. It appears that short bursts are scattered significantly on their hosts, occurring in typically fainter regions, and at larger offsets than their long cousins \\citep{LGRBs_in_bright_regions,Fong_2010_SGRBhosts_HST,Church_NS_offsets}. These properties can naturally be explained if the progenitor is a merger of two compact objects (e.g., neutron star -- neutron star (NS), collapsing to a black hole by accretion-induced collapse (AIC), or neutron star -- black hole (BH)) \\citep{mergingNSbinary_could_emit_gammarays,FryerNS_dist,spatialdistns_ns,Fong_2010_SGRBhosts_HST}. This model has proved extremely hard to test observationally, although the recent discovery of a possible kilonova in the SGRB 130603B offers support for the model for at least some SGRBs \\citep{tanvir13,berger13}.  A key distinguishing factor between different intrinsically ancient progenitor populations, such as accretion-induced collapse (AIC) of white dwarfs to neutron stars \\citep{Other_SGRB_progenitors,Metzger_2006_Magnetars}, and compact binary mergers, comes from the dynamics of the systems themselves. In a double compact object (DCO) binary, a combination of natal kicks, and mass loss from the binary at the time of each supernova, can act to provide % the systems with space velocities of several hundred $\\mathrm{km}\\s^{-1}$ (e.g. \\citealt{2010_8_ns_binaries}). Integrated over the lifetime of the binary of $10^7-10^{10}$ years this can correspond to distances of tens of $\\mathrm{kpc}$ from their birth sites, although the fraction of binaries that attain these distances remains uncertain \\citep{Belczynski_2006_NSmerger,Church_NS_offsets}. For extremely high kicks, or relatively low mass host galaxies, the binary may escape the galactic potential of its host altogether. Hence, a population of SGRBs in intergalactic space would offer strong support for a binary merger model for their progenitors \\citep{Bloom_2006_diverse_progenitors,Berger_2010_nohost}.  However, determining the offset of a burst from its host is non-trivial when there is no obvious parent galaxy coincident with or close to the burst location, i.e. where there is no outstanding candidate with an extremely low probability of chance alignment. Probabalistic methods based on the sky density of galaxies are often used to argue for a host association (e.g. \\citealt{GRBoffset, Bloom_2007_060502B_putative_hosts, AL_random_associations, Berger_2010_nohost}). However, the sky density is such that for that any random position on the sky is likely to be within a few arcseconds of a galaxy with $R<25$. In other words, in many cases we cannot \\emph{strongly} identify the host galaxy (probability of $\\lesssim1\\%$), % which can lead to mis-identifications. % A second problem is that to the limits of our ground-based (or even {\\em Hubble Space Telescope}) observations, we probe a reducing fraction of the galaxy luminosity function as we move to higher redshift. Hence there is a potential for confusion between GRBs which have been kicked far, and hence are well offset, % from relatively local hosts and those which lie within fainter galaxies at high redshifts. This problem is particularly acute if the high redshift galaxies host primarily old stellar populations, and hence exhibit only weak rest frame UV (observer-frame optical) emission.   In this sense the ``hostless\" problem for SGRBs in not that there are a lack of candidate hosts; in all lines of sight there will be plausible parent galaxies within a few tens of kpc in projection. Frequently the probability of chance alignment with at least one of these is small, and may be suggestive of kicks to the SGRB progenitors \\citep{Berger_2010_nohost}. Instead, the problem in these hostless cases is the difficulty in determining uniquely the parent galaxy % (from e.g. several with similar probabilities, or underlying larger scale structure). This means we are unable to make full use of the diagnostic information contained in the offset distribution and the properties of the hosts for improving our knowledge of SGRB progenitors.  Obtaining the redshift for the GRB using the afterglow would allow us to narrow our search to hosts within a small redshift range. However, for short bursts the faintness of their afterglows \\citep{Kann_2011_OAG_TypeI_vs_TypeII} means that redshifts are difficult to obtain in practice, and in nearly all cases to-date redshifts for SGRBs have been inferred from their presumed host galaxy rather than from the burst itself (e.g. \\citealt{Hjorth_2005_050709_OA,Berger_2005_050724_host,Rowlinson_2010_080905A_Opt,Fong_2011_100117A_host}). In fact, % even accounting for the faint continuum, in some cases the lack of any absorption features  % also gives an indication that the burst is not in a dense interstellar medium (ISM) \\citep{Piranomonte_2008_070707_Beta-OX,Berger_GCN_100117A_Gemini_spec}. %   Based on the optically localised SGRB sample outlined in section \\ref{SGRB_sample} with details in appendix \\ref{SGRB_sample_table}, $\\sim70\\%$ % have apparently well-associated host galaxies, while the remainder are apparently hostless. This may offer evidence for kicks. Here, we present the discovery and subsequent observations of the optical afterglows and host limits of a further two hostless bursts, GRB\\,090305A (see also \\citealt{Berger_2010_nohost} and \\citealt{NicuesaGuelbenzu_2012_preprint_090305}), GRB\\,091109B and deep limits to a third, GRB\\,111020A (see also \\citealt{Fong_GCN_111020A_Gemini}). These GRBs are all unambiguously of the short-hard class, % with $T_{90}<0.5\\s$ and prompt emission which is spectrally hard. Indeed, even given the recent work suggesting a shorter dividing line between the short- and long-GRB populations \\citep{Bromberg_2013_short_swift}, these bursts would remain in the short-hard class. We consider the extent to which our current detection limits probe the galaxy luminosity function as a function of redshift, and what this implies for hostless GRBs more generally.  These hostless GRBs could reside within relatively high redshift (but so far unseen) host galaxies or have travelled far from their low redshift ($z<1$) hosts, perhaps within the intergalactic medium (IGM). Clues to their origins may come from studies of the most likely hosts amongst the nearby galaxies on the sky. We present an alternative diagnostic tool developed by taking random positions and comparing them to the distribution of galaxies on the sky, thus reproducing the sort of analysis performed when looking for a short GRB host. From this we directly determine the chance probability of association ($P_{\\rm chance}$) and a radius within which we can confidently state a host association. Doing this avoids assumptions as to the functional form of the number counts, and naturally encapsulates field to field variance and clustering, resulting in more robust assessments of the chance probability than the traditional route of using number counts alone.  % ",
        "conclusions": "We have looked at the afterglow properties of three apparently hostless SGRBs: GRB\\,090305A, GRB\\,091109B and GRB\\,111020A. The former, in particular, had a very faint X-ray counterpart, only identified due to the detection of the optical afterglow within the BAT error circle. Detection of the afterglow of GRB\\,091109B in the $R$ band and GRB\\,090305A in the $g^{\\prime}$ band allows us to place upper redshift limits of $z\\lesssim5$ and $z\\lesssim3.5$, respectively.  Deep optical observations at the GRB position after the afterglow had faded allows us to put constraints on any coincident host galaxy, specifically the $3\\sigma$ limiting magnitude for GRB\\,090305A is $r>25.69$, for GRB\\,091109B the limits are $R>25.80$, $K>22.23$ and $J>21.99$ and for GRB\\,111020A is $R>24.23$ (a deeper limit of $i^{\\prime}>24.4$ has also be placed by \\citealt{Fong_2012_111020A_preprint_Xray}). Although $r^{\\prime}$ band observations make elliptical galaxies in particular difficult to detect with lower redshift limits, use of a deep limit in the $K$ band would even the chances of detecting any type of host galaxy. % Using the deepest limiting magnitudes for GRBs 090305A and 091109B we find that out to $z=1$ we uncover $\\sim75\\%$ of the integrated galaxy luminosity for an Sbc type galaxy, $\\sim85\\%$ for an Irr galaxy and $55\\%$ for an Ell galaxy. GRB\\,111020A which, even using the deepest limit available from \\citealt{Fong_2012_111020A_preprint_Xray}, has a shallower limit, would uncover $\\sim50\\%$ (Sbc galaxy). It is unclear as to the status of GRB\\,111020A as a hostless burst due to the presence of an unresolved object $0.7\\,''$ from its position which would have $P_{\\rm chance}<0.01$\\footnote{This case is also complicated by the relatively low Galactic latitude, which means that foreground stars significantly outnumber background galaxies in our images.}. However, for other GRBs considered here as well as no coincident host detection, we also find no host that can be confidently identified using $P_{\\rm chance}$ values ($<1\\%$). %  These GRBs represent a growing population of optically localised SGRBs with no obvious host galaxy. We have considered two possible origins for these hostless SGRBs. The first is that they originate from a higher redshift, and so far unseen population of SGRBs, while the second is that they lie at lower redshift, and are kicked from local, and relatively bright host galaxies.  To address these issues we developed a diagnostic to assess the significance of the association of any given galaxies with a SGRB, and compared the properties of the sample of bursts with those of random positions on the sky. These results suggest that hostless SGRBs as a population have a correlation with structure at small angular scales, more so that ``average\" random lines of sight. This perhaps offers evidence that the hostless SGRBs are in fact associated either with these bright galaxies, or with fainter galaxies associated with the same large scale structure. In this case the offsets are either as a result of large scale natal kicks to the progenitors, or of their dynamical formation within globular clusters.   We note that a similar study by \\cite{Berger_2010_nohost} concluded that large offsets of $15-70\\,\\mathrm{kpc}$ from relatively low redshift galaxies are their preferred explanation. They found for a high redshift solution the constraints on any underlying host galaxy implied a bimodal population of SGRBs with peaks at $z\\sim0.5$ and $z\\sim3$. They also allow for a minor contribution from NS binary mergers in globular clusters. % Our work broadly agrees with the results considered by  \\cite{Berger_2010_nohost}, although critically extends it to consider the true distribution of galaxies on the sky (rather than average number counts), utilizing this comparison to make statistical statements on the population as a whole.  Ultimately, if hostless GRBs are present at low redshift deeper observations of their locations will continue to yield null detections of their host galaxies. However, for bursts associated with structure at lower redshift we may be able to ascertain if they are hosted within globular clusters via deep observations with either the Hubble Space Telescope (to $z<0.1$) or the James Webb Space Telescope (to $z<0.2$). Such a detection could offer strong evidence for the origin of SGRBs in compact binary mergers."
    },
    "1402/1402.0882_arXiv.txt": {
        "abstract": "The first generation of stars produces a background of Lyman-Werner (LW) radiation which can photo-dissociate molecular hydrogen, increasing the mass of dark matter halos required to host star formation.  Previous studies have determined the critical mass required for efficient molecular cooling with a constant LW background.  However, the true background is expected to increase rapidly at early times.  Neglecting this evolution could underestimate star formation in small halos that may have started to cool in the past when the LW intensity was much lower.   Background evolution is a large source of uncertainty in pre-reionization predictions of the cosmological 21cm signal, which can be observed with future radio telescopes.  To address this, we perform zero-dimentional one-zone calculations that follow the density, chemical abundances, and temperature of gas in the central regions of dark matter halos, including hierarchical growth and an evolving LW background.  We begin by studying the physics of halos subjected to a background that increases exponentially with redshift.  We find that when the intensity increases more slowly than $J_{\\rm LW}(z) \\propto 10^{-z/5}$, cooling in the past is a relatively small effect.  We then self-consistently compute the cosmological LW background over $z=15-50$ and find that cooling in the past due to an evolving background has a modest impact.  Finally, we compare these results to three-dimensional hydrodynamical cosmological simulations with varying LW histories.  While only a small number of halos were simulated, the results are consistent with our one-zone calculations. ",
        "introduction": "The first stars in the universe formed during the period between recombination ($z\\sim1100$) and cosmic reionization ($z\\sim7-15$).  Fortunately, due to the lack of heavy elements during that epoch, modeling these stars is a relatively clean theoretical problem.  Primordial, metal-free (Pop III) stars are believed to form in small, $\\sim 10^5 M_\\odot$ dark matter halos that undergo efficient molecular hydrogen cooling  \\citep{2002Sci...295...93A,1996ApJ...464..523H,1997ApJ...474....1T,2001ApJ...548..509M}.  While it will not be possible in the near future to detect individual objects at these extremely high redshifts, future measurements of the 21cm hydrogen line will have the ability to test theoretical predictions \\citep[for reviews see][]{2006PhR...433..181F, 2012RPPh...75h6901P}.  Future observatories such as the Square Kilometer Array (SKA) \\citep{2004NewAR..48.1029C}, sensitive to redshifts as high as $z \\sim 30$, could infer the properties of early stars through their impact on the intergalactic medium (IGM).  An important aspect of Pop III star formation is negative feedback from the Lyman-Werner (LW) background  \\citep{1997ApJ...476..458H}.  LW photons (11.2--13.6 eV) destroy molecular hydrogen and raise the minimum mass of dark matter halos where star formation can occur.  Previous works utilizing both hydrodynamical simulations \\citep[e.g.][]{2001ApJ...548..509M,2007ApJ...671.1559W,2008ApJ...673...14O} and analytic calculations \\citep[e.g.][]{1997ApJ...474....1T,2000ApJ...534...11H} have estimated the critical dark matter halo mass required for efficient molecular cooling as a function of background intensity and redshift, $M_{\\rm crit}(J_{\\rm LW},z)$.    To compute $M_{\\rm crit}(J_{\\rm LW},z)$, all of these studies made the simplifying assumption of a constant LW background intensity.  However, the strength of the background is expected to increase rapidly as a function of time.  Ignoring this time dependence could potentially underestimate the amount of star formation.  This is because a halo that starts to cool in the past, when the background was lower, may continue to cool even if the background increases such that its mass is below the critical value at a later time.  The impact of time evolution in the LW background on early star formation is currently a large source of uncertainty in predicting the pre-reionization 21cm signal \\citep{2013MNRAS.432.2909F,2012Natur.487...70V, 2012ApJ...760....3M}.  \\cite{2013MNRAS.432.2909F} account for this by assuming that $M_{\\rm crit}(J_{\\rm LW},z)$ should be evaluated with the value of $J_{\\rm LW}$ sometime in the past, when irreversible cooling presumably began.  A range of values for this time delay were selected based on the spherical collapse model, leading to significant changes in the 21cm power spectrum.  The appropriate choice of this delay is unclear \\emph{ab initio}.  In this paper, we perform zero-dimensional one-zone calculations of negative LW feedback with a time-evolving background in an effort to better understand the physics of primordial star formation.  Our calculations follow the density and chemistry of star-forming dark matter halos that grow hierarchically within a rapidly increasing LW background.  To track the growth of dark matter halos we generate Monte Carlo (MC) merger trees and follow the mass of the main progenitor (found by starting at the final halo and going back in time along the branch with the most-massive progenitor at each merger).  Our analysis of these one-zone calculations is divided into two main parts.  First, we consider the impact of a time-dependent LW background on star formation in individual dark matter halos.  For a given halo mass, if and when cooling occurs depends on the particular merger history of each halo.  For this reason, rather than $M_{\\rm crit}(J_{\\rm LW},z)$, it is more precise to describe the cooling threshold in terms of the fraction of halos at a given mass and redshift that have cooled, $f_{\\rm cool}(z,M,J_{\\rm LW}(z))$.  Here $J_{\\rm LW}(z)$ represents the history of the background at all redshifts. We compute $f_{\\rm cool}$ at $z=25$ for a range of exponentially increasing $J_{\\rm LW}(z)$.  Cooling in the past due to a lower LW background is a relatively small effect for backgrounds that increase more slowly than $J_{\\rm LW} \\propto 10^{-z/5}$.  In the second portion of our analysis, we self-consistently compute the cosmological $J_{\\rm LW}(z)$ across all relevant redshifts.  Because the LW background increases sufficiently slowly, we find that cooling depends mainly on the current LW intensity and not on its history.  We compare these results to three-dimensional cosmological simulations, including hydrodynamics and primordial chemistry, and find that they give the same main conclusion, the intensity of the LW background in the past does not have a strong impact on when a halo can start to cool.  This paper is structured as follows.  Is $\\S2$ we analyze the impact of a time-evolving LW background on individual dark matter halos.  We describe our computational method and prescription for generating MC merger trees in the first two subsections.   In $\\S$2.3, we demonstrate how $f_{\\rm cool}$ changes for different LW background histories.  In $\\S$3, we use our method to self-consistently compute the background strength over all relevant redshifts.  In $\\S$4, we present results from hydrodynamic cosmological simulations with evolving LW backgrounds.  We discuss our results and conclusions in $\\S$5.  Except for the three-dimensional numerical simulations, we assume a $\\Lambda$CDM cosmology consistent with the latest constraints from Planck \\citep{2013arXiv1303.5076P}: $\\Omega_\\Lambda=0.68$, $\\Omega_{\\rm m}=0.32$, $\\Omega_{\\rm b}=0.049$, $h=0.67$, $\\sigma_8=0.83$, and $n_{\\rm s} = 0.96$.  The slightly different cosmology used in the numerical simulations should not have a significant impact on our results. ",
        "conclusions": "Negative feedback from the LW background can destroy molecular hydrogen and prevent star formation in the high-redshift universe.  Previous studies calculating the critical dark matter halo mass required to host star formation have all considered a constant LW background.  This can underestimate the amount of star formation if cooling is able to occur in the past when the background was much lower.  This effect has been a major source of uncertainty in semi-numerical predictions of the pre-reionization 21cm signal \\citep{2013MNRAS.432.2909F,2012Natur.487...70V,2013arXiv1306.2354F}.  We have completed the first comprehensive analysis of negative LW feedback that considers a time-dependent background.  We performed one-zone calculations that follow the density, temperature, and chemistry of gas in the central regions of dark matter halos subjected to a rapidly increasing LW background.  To characterize the mass threshold for efficient molecular cooling we computed $f_{\\rm cool}(z,M,J_{\\rm LW}(z))$, the fraction of halos that cool sometime before redshift $z$.  In $\\S$2, we determined $f_{\\rm cool}$ for a range of artificial LW background histories that grow exponentially according to Eqn.~\\ref{test_bg}.  We found that for backgrounds which increase more slowly than $\\alpha_{\\rm LW} \\sim 5$, the effect of time dependence is relatively small (see Fig.~\\ref{cool_frac}).  In \\S3, we self-consistently computed the cosmological background from $z=15-50$ including the effect of time-dependence.  We compared this with the background obtained when $f_{\\rm cool}(z,M)$ is calculated assuming constant $J_{\\rm LW}$ (set to the value at the previous redshift step).  We find that time-dependence has a relatively small impact ($<30$ percent).  We found similar results with both the $c_1=0.72$, $c_2=0.4$, and $c_3=0.25$ and $c_1=c_2=c_3=1$ parametrizations of our model.  Thus, even though our model does not exactly match the critical cooling mass taken from numerical simulations, our overall conclusion that mainly the current $J_{LW}$ and not its history impacts star formation, seems robust.  This is also supported by the cosmological simulations presented in \\S4.  Our results suggest that semi-numerical simulations of the large-scale distribution of the first stars will not be changed greatly by the LW time-dependance effect.  However, we point out that (even in the case of a constant background) $f_{\\rm cool}(M,z)$ provides a better description of the cooling threshold than a simple step function, because cooling depends not only on a halo's current state, but also on its individual detailed assembly history.  Using $f_{\\rm cool}$ calibrated with hydrodynamical simulations, rather than a cut-off at $M_{\\rm crit}$ could increase the accuracy of the predicted 21cm signal.  We have not performed a detailed comparison of our one-zone calculations and the numerical simulations.  In future work it will be beneficial to make detailed checks of our assumptions regarding the central density and cooling criterion. As our simulations followed only a limited number of halos, future larger simulations with more halos would help to confirm our results.  Finally, we note that throughout this analysis we have ignored the impact of large-scale baryon-dark matter streaming velocities \\citep{2010PhRvD..82h3520T,2012MNRAS.424.1335F}.  The $J_{\\rm LW}(z)$ we compute will only apply to large-scale ($\\sim 100 {\\rm Mpc}$) regions with low streaming velocity.  However, since the relative velocity effect tends to delay the star formation, it seems likely that our overall conclusions will apply to fast moving regions as well.  This should be checked in detail with future hydrodynamical simulations."
    },
    "1402/1402.5719_arXiv.txt": {
        "abstract": "\\noindent We derive the dynamics of the gravitational collapse of a homogeneous and spherically symmetric cloud in a classical set-up endowed with a topological sector of gravity and a non-minimal coupling to fermions. The effective theory consists of the Einstein-Hilbert action plus Dirac fermions interacting through a four-fermion vertex. At the classical level, we obtain the same picture that has been recently studied by some of us within a wide range of effective theories inspired by a super-renormalizable and asymptotically free theory of gravity. The classical singularity is replaced by a bounce, beyond which the cloud re-expands indefinitely. We thus show that, even at a classical level, if we allow for a non-minimal coupling of gravity to fermions, event horizons may never form for a suitable choice of some parameters of the theory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3415_arXiv.txt": {
        "abstract": "Considering six general relativistic, two-dimensional (2D) supernova (SN) explosion models of progenitor stars between 8.1 and 27\\,$M_\\odot$, we systematically analyze the properties of the neutrino emission from core collapse and bounce to the post-explosion phase. The models were computed with the \\textsc{Vertex-CoCoNuT} code, using three-flavor, energy-dependent neutrino transport in the ray-by-ray-plus approximation. Our results confirm the close similarity of the mean energies, $\\langle E\\rangle$, of $\\bar\\nu_e$ and heavy-lepton neutrinos and even their crossing during the accretion phase for stars with $M\\gtrsim 10\\,M_\\odot$ as observed in previous 1D and 2D simulations with state-of-the-art neutrino transport. We establish a roughly linear scaling of $\\langle E_{\\bar\\nu_e}\\rangle$ with the proto-neutron star (PNS) mass, which holds in time as well as for different progenitors.  Convection inside the PNS affects the neutrino emission on the 10--20\\% level, and accretion continuing beyond the onset of the explosion prevents the abrupt drop of the neutrino luminosities seen in artificially exploded 1D models.  We demonstrate that a wavelet-based time-frequency analysis of SN neutrino signals in IceCube will offer sensitive diagnostics for the SN core dynamics up to at least $\\sim$10\\,kpc distance. Strong, narrow-band signal modulations indicate quasi-periodic shock sloshing motions due to the standing accretion shock instability (SASI), and the frequency evolution of such ``SASI neutrino chirps'' reveals shock expansion or contraction. The onset of the explosion is accompanied by a shift of the modulation frequency below 40--50\\,Hz, and post-explosion, episodic accretion downflows will be signaled by activity intervals stretching over an extended frequency range in the wavelet spectrogram. ",
        "introduction": "\\label{sec:intro}  Neutrinos play a fundamental role in core-collapse supernovae: Not only do they carry away most (several $10^{53} \\ \\mathrm{erg}$) of the gravitational binding energy liberated during the collapse of the inner shells of the progenitor to a proto-neutron star, they are also the driving agent for the supernova explosion in the most popular scenario for shock revival, the ``delayed neutrino-driven mechanism'' of \\citet{bethe_85}; see \\citet{janka_12,janka_12b,burrows_13,kotake_12b} for current reviews. Once the explosion has been successfully initiated and post-shock accretion has ceased, the continuing neutrino emission from the proto-neutron star drives a baryonic wind that has long been considered as a site of interesting nucleosynthesis (see, e.g., \\citealt{arcones_13}).  Moreover, the observation of supernova neutrinos from a future Galactic event could provide tremendous insights into the dynamics deep inside the stellar core and serve to constrain unknown particle physics.  The detection of some two dozen neutrinos from Supernova SN~1987A \\citep{hirata_87,bionta_87,alekseev_87} already confirmed the basic picture of neutrino emission in core-collapse supernovae, suggesting the emission of $\\sim 3 \\times 10^{53} \\ \\mathrm{erg}$ with a time-averaged neutrino temperature of $\\sim 4 \\ \\mathrm{MeV}$ from a neutrinosphere of the order of a few $10 \\ \\mathrm{km}$ with a total signal duration of a few seconds (see, e.g., \\citealt{arnett_89,burrows_90,koshiba_92} for an overview). The signal from SN~1987A also provided constraints on the mass, electric charge, magnetic moment and lifetime of neutrinos as well as indirect constraints on the mass of a hypothetical axion (see \\citealt{burrows_90,keil_97} and references therein).  Due to much higher event rates, present and next-generation detectors will allow for the reconstruction of the time-dependent neutrino signal in much greater detail (see, e.g., \\citealt{abe_12,wurm_12} and, for a general overview, \\citealt{scholberg_12}), which could allow far-reaching conclusions both on the dynamics in the supernova core and open questions in neutrino physics such as the mass hierarchy \\citep{dighe_00,serpico_12}.  Quantitatively accurate predictions for the neutrino emission from supernova simulations are an indispensable prerequisite for properly interpreting the neutrino signal from a prospective Galactic event. Nowadays, sophisticated methods for the solution of the neutrino transport problem are available for this purpose. The most advanced schemes for neutrino transport in \\emph{spherically symmetric} (1D) neutrino hydrodynamics simulations rely either on the direct solution of the energy-dependent general relativistic \\citep{yamada_97,liebendoerfer_04} or Newtonian \\citep{mezzacappa_93} Boltzmann equation, or on a variable Eddington factor method with Boltzmann closure (Newtonian: \\citealt{burrows_00}, pseudo-relativistic: \\citealt{rampp_02}, general relativistic: \\citealt{mueller_10,roberts_12a}).  These methods have been used to address different phases of the neutrino emission from core-collapse supernovae.  Several studies focused specifically on the neutronization burst \\citep{kachelriess_05,langanke_08} and the rise of the electron antineutrino and the heavy flavor neutrino luminosity \\citep{serpico_12} around shock breakout (with an emphasis on the detectable signal). The dependence of the neutrino emission during the pre-explosion (accretion) phase on the progenitor properties and the equation of state (EoS) of nuclear matter has been the subject of a larger number of papers \\citep{thompson_03,liebendoerfer_03,sumiyoshi_05,buras_06_b,fischer_09,oconnor_13,nakazato_13}. \\citet{sumiyoshi_08} and \\citet{fischer_09} also cover the case of black hole formation in failed supernovae.  \\citet{lentz_12a,lentz_12b} recently studied the impact of variations in the microphysics and approximations in the neutrino transport sector and emphasized the importance of a rigorous treatment of general relativity, observer corrections, and energy-exchanges in scattering reactions on electrons and nucleons for accurate predictions of neutrino luminosities and mean energies.  1D simulations using multi-group Boltzmann or variable Eddington factor transport have likewise been used to predict the signal from the Kelvin-Helmholtz cooling phase after the successful initiation of an explosion \\citep{huedepohl_10,fischer_10,fischer_12,roberts_12a,roberts_12b,martinez_12}. However, with the exception of electron-capture supernovae \\citep{huedepohl_10,fischer_10}, predictions for the neutrino signal from the post-explosion phase are currently based on \\emph{artificial} explosion models in which shock revival is achieved by manually boosting neutrino heating in the gain layer, or even start from initial models of the proto-neutron star constructed by hand \\citep{roberts_12a,roberts_12b}.  According to our current understanding, neutrino-driven core-collapse supernovae are inherently \\emph{multi-dimensional} (multi-D).  Hydrodynamic instabilities like buoyancy-driven convection in the neutrino-heated gain region \\citep{bethe_90,herant_92,herant_94,burrows_95,janka_96,mueller_97} and the standing accretion shock instability (``SASI'', \\citealp{blondin_03,blondin_06,foglizzo_06,ohnishi_06, foglizzo_07,scheck_08,iwakami_08,iwakami_09,fernandez_09,fernandez_10}) were found to play a crucial role in the explosion mechanism, and they are no less relevant for the neutrino emission. However, predictions for the neutrino signal from multi-D supernova simulations with a transport treatment on par with the best available 1D models are still scarce. Attempts at a rigorous solution of the 3D Boltzmann equation \\citep{kotake_12,radice_13,peres_13} are yet in their infancy.  The best available studies of multi-D effects on the neutrino emission therefore either rely on Newtonian 2D multi-angle transport with considerable compromises in the microphysics and omission of energy-bin coupling \\citep{ott_08_a,brandt_10} or on multi-group variable Eddington factor transport using the ``ray-by-ray-plus approximation'' of \\citet{buras_06_a,bruenn_06} in 2D \\citep{marek_08,lund_10} and 3D \\citep{tamborra_13}. Predictions for the explosion phase are currently available only from parameterized models with gray transport and an excised neutron star core \\citep{mueller_e_12,lund_12}. Despite the different methodologies there is a consensus that the neutrino signal is considerably affected by multi-D instabilities. Asymmetric accretion onto the proto-neutron star gives rise to spatial anisotropies and temporal variations in the neutrino emission. These spatio-temporal variations show a distinct imprint of the different hydrodynamic instabilities in the supernova core, with SASI oscillations leading to fluctuations in the neutrino signal with a rather well-defined frequency around $100 \\, \\mathrm{Hz}$, and convective overturn resulting in more stochastic variations \\citep{tamborra_13}. After the onset of the explosion, the neutrino emission is typically characterized by large-scale spatial anisotropies as the accretion flow onto the proto-neutron star (which usually subsists for several hundreds of milliseconds after shock revival) becomes highly asymmetric.   In this paper, we reexamine the neutrino emission in core-collapse supernovae on the basis of state-of-the-art axisymmetric (2D) general relativistic simulations with energy-dependent three-flavor neutrino transport. We consider a wide range of progenitors from the lowest-mass iron cores, where multi-D effects play a minor role for the neutrino emission, through more massive models with vigorous convection and SASI. Different from previous studies of neutrino emission in multi-D simulations of the post-bounce phase \\citep{ott_08_a,marek_08,lund_10,brandt_10,tamborra_13}, we explore both the pre-explosion and the explosion phase. We analyze both the overall secular evolution of the neutrino emission as well as spatio-temporal variations on shorter time-scales. In order to connect with future neutrino observations, we show how the quantitative analysis of the signal from a prospective Galactic event in a detector like IceCube \\citep{abbasi_11,salathe_12} could provide detailed time-dependent information about the dynamics in the supernova core.   Our paper is structured as follows: In Section~\\ref{sec:numerics}, we outline the numerical methods and the input physics of our simulations, including the progenitor models. In Section~\\ref{sec:overview}, we review the secular evolution of the spherically averaged neutrino luminosities and mean energies.  In Section~\\ref{sec:variations}, we discuss the spatio-temporal variations in the neutrino emission on the basis of simulated IceCube signals. By means of a time-frequency analysis using wavelet transforms, we show that the IceCube signal can reveal the detailed time-dependence of the SASI sloshing frequency, the rough evolution of the shock radius, the onset of the explosion, and early fallback onto the proto-neutron star through newly-formed accretion funnels. In Section~\\ref{sec:summary}, we summarize our findings, examine their robustness, and discuss further implications for the observation of a future Galactic supernova in gravitational waves and neutrinos.   \\begin{table*} \\caption{Neutrino physics input \\label{tab:rates} } \\begin{center} \\begin{tabular}{cc} \\hline \\hline process  & reference \\\\ \\hline $\\nu A \\rightleftharpoons \\nu A$ & \\citet{horowitz_97} (ion-ion correlations) \\\\ & \\citet{langanke_08} (inelastic contribution) \\\\ $\\nu \\ e^\\pm \\rightleftharpoons \\nu \\ e^\\pm$  &  \\citet{mezzacappa_93} \\\\ $\\nu \\ N \\rightleftharpoons \\nu \\ N$          & \\citet{burrows_98}\\tablenotemark{a} \\\\ $\\nu_e\\ n \\rightleftharpoons e^-\\ p$      & \\citet{burrows_98}\\tablenotemark{a} \\\\ $\\bar{\\nu}_e \\ p \\rightleftharpoons e^+ \\ n$  & \\citet{burrows_98}\\tablenotemark{a} \\\\ $\\nu_e \\ A' \\rightleftharpoons e^- \\ A$  & \\citet{langanke_03} \\\\ $\\nu\\bar{\\nu} \\rightleftharpoons e^- \\ e^+ $ & \\citet{bruenn_85,pons_98} \\\\ $\\nu\\bar{\\nu} \\ NN \\rightleftharpoons NN $ & \\citet{hannestad_98} \\\\ $\\nu_{\\mu,\\tau}\\bar{\\nu}_{\\mu,\\tau} \\rightleftharpoons \\nu_{e}\\bar{\\nu}_{e} $ & \\citet{buras_03} \\\\ $\\stackrel{(-)}{\\nu}_{\\mu,\\tau}\\stackrel{(-)}{\\nu}_e \\rightleftharpoons \\stackrel{(-)}{\\nu}_{\\mu,\\tau}\\stackrel{(-)}{\\nu}_e$ & \\citet{buras_03} \\\\ \\hline \\hline \\end{tabular} \\tablenotetext{1}{ Note that these reaction rates account for nucleon thermal motions, phase-space blocking, energy transfer to the nucleon associated with recoil (``non-conservative'' or ``non-isoenergetic'' scattering), and nucleon correlations at high densities.  Moreover, we include the quenching of the axial-vector coupling at high densities \\citep{carter_02}, correction to the effective nucleon mass \\citep{reddy_99}, and weak magnetism effects \\citep{horowitz_02}. However, we ignore nucleon potential effects \\citep{martinez_12,roberts_12c}, which are of minor importance during the accretion phase \\citep{martinez_12}.} \\end{center} \\end{table*}   \\begin{table*} \\caption{Model setup \\label{tab:model_setup} } \\begin{center} \\begin{tabular}{cccccccc} \\hline \\hline & metallicity & angular    & explosion  & time of                   &      & progenitor & simulation\\\\ progenitor  & $Z/Z_\\odot$ & resolution & obtained   & explosion\\tablenotemark{a} & EoS & reference  & reference\\\\ \\hline u8.1   & $10^{-4}$ & $1.4^\\circ$ & yes        & $175 \\ \\mathrm{ms}$ & LS180  & A.~Heger (private communication) & \\citet{mueller_12b} \\\\ z9.6    & $0$       & $1.4^\\circ$ & yes        & $125 \\ \\mathrm{ms}$ & LS220  & A.~Heger (private communication) & \\citet{mueller_13}\\\\ s11.2   & $1$       & $2.8^\\circ$ & yes        & $213 \\ \\mathrm{ms}$ & LS180  & \\citet{woosley_02} & \\citet{mueller_12a}\\\\ s15s7b2 & $1$       & $2.8^\\circ$ & yes        & $569 \\ \\mathrm{ms}$ & LS180  & \\citet{woosley_95} & \\citet{mueller_12a} \\\\ s25   & $1$       & $1.4^\\circ$ & no         & ---                 & LS220  & \\citet{woosley_02} & \\citet{mueller_13} \\\\ s27   & $1$       & $1.4^\\circ$ & yes        & $209 \\ \\mathrm{ms}$ & LS220  & \\citet{woosley_02} & \\citet{mueller_12b} \\\\ \\hline \\hline \\end{tabular} \\tablenotetext{1}{Defined as the point in time when the average shock radius $\\langle r_\\mathrm{sh} \\rangle$ reaches $400 \\ \\mathrm{km}$.} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:summary} Based on six general relativistic 2D simulations of progenitors between $8.1 M_\\odot$ and $27 M_\\odot$, we presented a detailed analysis of the neutrino emission from core-collapse supernovae from the bounce to the explosion phase. We discussed both the secular evolution of the total, angle-integrated neutrino emission, which is largely regulated by the continuous accretion of material and the contraction of the proto-neutron star, as well as spatio-temporal variations in the neutrino emission caused by nonaspherical hydrodynamic instabilities like convection and the SASI. Using simulated signals of a future Galactic supernova in IceCube \\citep{abbasi_11,salathe_12}, we studied the observability of these spatio-temporal variations and showed how detailed information about the supernova core could be extracted from the observed neutrino signal by means of a time-frequency analysis.  Our most salient findings can be summarized as follows: \\begin{enumerate} \\item Prior to the onset of the explosion, the evolution of the total neutrino flux and the neutrino mean energies is in qualitative agreement with recent 1D models of the accretion phase \\citep{liebendoerfer_03,liebendoerfer_04,buras_06_b,marek_09,fischer_09,oconnor_13} although proto-neutron star convection affects the neutrino emission on the level of $10 \\ldots 20 \\%$ \\citep{buras_06_b}. Similar to other modern neutrino hydrodynamics simulations, and in stark contrast to some earlier 1D models from the 1980's and 1990's, our 2D models are characterized by very similar mean energies of electron antineutrinos and $\\mu/\\tau$ neutrinos.  In all but the least massive progenitors ($8.1 M_\\odot$ and $9.6 M_\\odot$), we even observe a crossing of the mean energies already within the first few hundreds of milliseconds of post-bounce accretion as a result of a temperature inversion near the neutrinosphere. \\item The mass of the proto-neutron star emerges as the single most important parameter regulating the secular rise of the mean neutrino energies, at least for a given equation of state. For individual models, we find that the electron antineutrino mean energy $\\langle E_{\\bar{\\nu}_e} \\rangle$ scales very well with the proto-neutron star mass $M$, \\begin{equation} \\langle E_{\\bar{\\nu}_e}\\rangle \\propto M. \\end{equation} The proportionality constant is slightly progenitor-dependent. This scaling is a consequence of the roughly adiabatic stratification of the hot accretion mantle of the proto-neutron star between nuclear density and the neutrinosphere and the steepness of the mass-radius relation between $0.5 M_\\odot$ and $2 M_\\odot$ for nucleonic equations of state \\citep{lattimer_07,steiner_10}. \\item After shock revival, our 2D models differ considerably from artificial 1D explosion models \\citep{fischer_10}. We find that there is no abrupt drop of the electron neutrino and antineutrino luminosity to indicate the onset of the explosion because accretion downflows persist for a long time and transport fresh material to the cooling region.  The accretion luminosity therefore remains high and can even rise transiently due to early fallback. \\item The neutrino signal from a Galactic supernova could provide \\emph{time-dependent} information about the dynamics in the supernova core at least for distances $\\lesssim 10 \\, \\mathrm{kpc}$ if the survival probability of electron antineutrinos is high. In a detector like IceCube, strong SASI activity will reveal itself by a strong, narrow-banded modulation of the detection rate with a period of $T_\\mathrm{SASI}$ that is directly related to the average shock radius $r_\\mathrm{sh}$ and the proto-neutron star radius $r_\\mathrm{PNS}$: \\begin{equation} T_\\mathrm{SASI} \\approx 19 \\, \\mathrm{ms} \\left(\\frac{r_\\mathrm{sh}}{100 \\, \\mathrm{km}}\\right)^{3/2} \\ln \\left(\\frac{r_\\mathrm{sh}}{r_\\mathrm{PNS}} \\right). \\end{equation} The amplitude of the signal modulation will depend on the observer direction, as will the presence of an overtone at $T_\\mathrm{SASI}/2$. For sustained SASI activity, the observed ``SASI chirp'' signal will directly reveal the expansion or contraction of the shock. Based on model s25, one might speculate that failing supernovae with a sufficient delay time to black hole collapse will generally be accompanied by a such ``SASI neutrino chirp''. \\item A wavelet-based time-frequency analysis of the observed signal in IceCube could also help to pinpoint the onset of the explosion. We predict that the explosion will be accompanied by a shift of the typical modulation frequency of the signal below $40 \\ldots 50 \\, \\mathrm{Hz}$. Moreover, the early fallback of shocked material onto the proto-neutron star through new downflows in (initially) weak explosions will lead to a detectable burst-like \\emph{increase} in the electron antineutrino luminosity. In the wavelet spectrogram of the signal, such events would manifest themselves as localized vertical stripes stretching over an extended frequency range. \\end{enumerate}    The temporal variations in the neutrino signal are obviously an intriguing diagnostics for the dynamics in the supernova core, but their robustness could be a potential concern. Our results are based on axisymmetric models (instead of full 3D simulations), rely on the assumption of a normal neutrino mass hierarchy and disregard possible non-linear flavor conversion due to neutrino-neutrino refraction.  In our opinion, it is mainly the uncertainties in the neutrino physics that make it difficult to decide about the detectability of spatio-temporal variations in the neutrino signal. In the worst case of an inverted mass hierarchy, MSW conversion in the stellar envelope could lead to a complete swap of $\\bar{\\nu}_e$ and $\\bar{\\nu}_\\mu$ \\citep{dighe_00}. As the contribution of the accretion luminosity is small for $\\nu_\\mu$'s and $\\nu_\\tau$'s, this would reduce the amplitude of the signal fluctuations by a factor of several \\citep{tamborra_13}. Even under such pessimistic conditions, periodic fluctuations would remain detectable out to a few $\\mathrm{kpc}$ at least for models with strong SASI activity like s25. However, neither our most optimistic case nor this pessimistic scenario might be realized in nature. Neutrino-neutrino refraction effects could lead to flavor swap for at least for certain neutrino energies (see \\citealt{duan_09,duan_10} for a review) already close to the neutrinosphere, which would in turn affect the outcome of MSW flavor conversion in the envelope.  The precise conditions for non-linear neutrino flavor conversion (e.g.\\ the interplay of neutrino-neutrino refraction with the ordinary matter affect and the role of a non-axisymmetric neutrino distribution in phase space) are a matter of active debate \\citep{sawyer_09,sarikas_12a,sarikas_12b,saviano_12,cherry_12,mirizzi_12,raffelt_13a,raffelt_13b}. Moreover, the simple picture of adiabatic MSW flavor conversion in the envelope might not hold in the presence of sufficiently strong turbulent density perturbations in convection shells \\citep{kneller_10,kneller_13,lund_13}. These complications preclude any precise estimate about the amplitude of temporal fluctuations in the observed neutrino signal. However, uncertainties in the neutrino physics are unlikely to render these fluctuations completely undetectable for nearby supernovae driven by vigorous hydrodynamic iinstabilities.  Our use of the ray-by-ray-plus method for the neutrino transport instead of a full multi-angle treatment might also be a concern, but by reprocessing our ray-by-ray-plus results as described in Section~\\ref{sec:reconstruction}, we ensure that we do not grossly overestimate the amplitude of the fluctuating neutrino signal.   The restriction of our current models to 2D is probably less of a concern, especially for clearly SASI-dominated models like s25 and s27. As demonstrated by \\citet{hanke_13}, the SASI can grow no less vigorously in 3D under appropriate conditions, and the concomitant emission anisotropies are of similar magnitude as in our 2D models (see \\citealt{tamborra_13}). The viewing-angle dependence of SASI-induced neutrino flux variations may of course be different in 3D, where the SASI can also develop a spiral mode, and where both the sloshing axis and the plane of the spiral mode can be time-dependent.  It is evident that the temporal variations of the observable neutrino signal could prove a powerful diagnostic in the case of Galactic supernova, revealing much more than the mere presence of the SASI and an ``average SASI'' frequency as discussed in previous studies \\citep{marek_09,ott_08_a,lund_10,brandt_10,lund_12,tamborra_13}. Together with time-dependent measurements of the neutrino flux and mean energy, the temporal variations in the neutrino signal could potentially provide much more than such overall constraints on the core mass and compactness as discussed previously in the literature \\citep{bruenn_87,burrows_88,oconnor_13}, especially if gravitational wave spectra were also available to determine the surface g-mode frequency $f$ of the proto-neutron star (with $f(t) \\propto G M/(r_\\mathrm{PNS}^2 \\langle E_{\\bar{\\nu}_e}\\rangle) \\left(1-G M/r_\\mathrm{PNS} c^2\\right)^{2}$ as shown in \\citealt{mueller_13}). The approximate mass-temperature relation $\\langle E_{\\bar{\\nu}_e}\\rangle \\propto M$ and the relation for the SASI chirp might then allow a (tentative) reconstruction of the parameters $M(t)$ (proto-neutron star mass), $r_\\mathrm{PNS}(t)$ (proto-neutron star radius), and $r_\\mathrm{sh} (t)$ (shock radius) of the accretion flow in the pre-explosion phase.  Given the manifold uncertainties in predicting precise neutrino signal templates (neutrino flavor conversion, opacities in dense matter, equation-of-state dependence), it is, of course, obvious that the ``signal inversion problem'' is bound to remain a serious challenge in neutrino astronomy that goes far beyond the scope of this paper.  We believe, however, that the work presented here may provide some useful ideas for interpreting future observations of core-collapse supernovae in neutrinos."
    },
    "1402/1402.3553_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small In this review we focus on the role jets and outflows play in the star and planet formation process. Our essential question can be posed as follows: are jets/outflows merely an epiphenomenon associated with star formation or do they play an important role in mediating the physics of assembling stars both individually and globally? We address this question by reviewing the current state of observations and their key points of contact with theory. Our review of jet/outflow phenomena is organized into three length-scale domains: Source and Disk Scales ($0.1-10^2$ au) where the connection with protostellar and disk evolution theories is paramount; Envelope Scales ($10^2-10^5$ au) where the chemistry and propagation shed further light on the jet launching process, its variability and its impact on the infalling envelope; Parent Cloud Scales ($10^5-10^6$ au) where global momentum injection into cluster/cloud environments become relevant. Issues of feedback are of particular importance on the smallest scales where planet formation regions in a disk may be impacted by the presence of disk winds, irradiation by jet shocks or shielding by the winds. Feedback on envelope scales may determine the final stellar mass (core-to-star efficiency) and envelope dissipation. Feedback also plays an important role on the larger scales with outflows contributing to turbulent support within clusters including alteration of cluster star formation efficiencies (feedback on larger scales currently appears unlikely). In describing these observations we also look to the future and consider the questions that new facilities such as ALMA and the Jansky Array can address. A particularly novel dimension of our review is that we consider results on jet dynamics from the emerging field of High Energy Density Laboratory Astrophysics (HEDLA). HEDLA is now providing direct insights into the 3-D dynamics of fully magnetized, hypersonic, radiative outflows. \\\\~\\\\~\\\\~} ",
        "introduction": "}  In many ways the discovery that star formation involves {\\it outflow} as well as {\\it inflow} from gravitational collapse marked the beginning of modern studies of the assembly of stars.  Jets and outflows were the first and most easily observed recognition that the narrative of star formation would include many players and processes beyond the spherical collapse of clouds.  The extraordinary progress made in the study of protostellar jets and outflows since the first discovery of Herbig-Haro (HH) objects (1950s), HH Jets (1980s) and molecular outflows (1980s) also reflects the growing power and sophistication of star formation science. The combination of ever higher resolution observational and computational methods, combined with innovative laboratory experiments, have allowed many aspects of the protostellar outflow problem to be clarified, though as we shall see crucial issues such as the launching process(es) remain debated.  Hypersonic collimated protostellar mass-loss appears to be a ubiquitous aspect of the star formation process. The observations currently indicate that most, if not all, low and high mass stars produce accretion-powered collimated ejections during their formation. These ejections are traced in two ways (see eg. the excellent PPV observational reviews by \\citet{2007prpl.conf..231R}, \\citet{2007prpl.conf..245A}, \\cite{2007prpl.conf..215B}).  First there are the narrow, highly-collimated ``jets'' of atomic and/or molecular gas with velocities of order $v \\sim 100-1000$ km/s ($v$ increasing with central source mass). These jets are believed to arise through magneto-hydrodynamic (MHD) processes in the rotating star-disk system. The other tracer are the less collimated more massive ``molecular outflows'' with velocities of order $v \\sim 1-30$ km/s which are believed to consist of shells of ambient gas swept-up by the jet bowshock and a surrounding slower wider-angle component. The fast and dense jet quickly escapes from the protostellar envelope (the still infalling remains of the original ``core'' from which the star formed) and propagates into the surrounding environment to become a ``parsec-scale outflow\". The less dense wide-angle wind and the swept-up outflow expand more slowly, carving out a cavity which widens over time into the envelope and the surrounding cloud. Most stars are born in clustered environments where the stellar separation is $< 1\\,pc$.  Thus, these large-scale outflows affect the interstellar medium within a cluster and, perhaps, the cloud as a whole. The important elements and processes on each scale from star to cloud are illustrated in Fig.\\,\\ref{Hartigan}.  Given their ubiquity and broad range of scales, a central question is whether jets and outflows constitute a mere epiphenomenon of star formation, or whether they are an essential component in the regulation of that process. In particular, winds/outflows are currently invoked to solve several major outstanding issues in star formation: (1) the low star formation efficiency in turbulent clouds (see eg. chapters by Padoan et al., Krumholtz et al.), (2) the systematic shift between the core mass function and the stellar initial mass function, suggesting a  core-to-star efficiency of only 30\\% (see eg. chapters by Offner et al., Padoan et al.), (3) the need to efficiently remove angular momentum from the young star and its disk. The former is important to avoid excessive spin-up by accretion and contraction, (cf chapter by Bouvier et al.) while the latter is required to maintain accretion at observed rates, in particular across the ``dead-zone\" where MHD turbulence is inefficient (cf. chapter by Turner et al.). Last but not least, protostellar winds may also affect disk evolution and planet formation through disk irradiation or shielding, and enhanced radial mixing of both gas and solids.  In this chapter we address this central question of outflow feedback on star and planet formation while also reviewing the current state of jet/outflow science. From the description above, it is clear that the degree of feedback will differ according to outflow properties on different spatial scales. The impact on the star and disk will depend on the physics of jet launching and angular momentum extraction (small scales). The impact on core-to-star efficiencies will depend both on the intrinsic jet structure and the jet propagation/interaction with surrounding gas. Finally, the impact on global star formation efficiency will depend on the overall momentum injection and on the efficacy of its coupling to cloud turbulence.  Thus in what follows we review the current state of understanding of protostellar jets and outflows by breaking the chapter into 3 sections through the following division of scales: Star and Disk  (1-10$^{\\rm 2}$\\,au); Envelope and Parent Clump (10$^{\\rm 2}$au - 0.5pc); Clusters and Molecular Clouds (0.5 - 10$^{\\rm 2}$\\,pc). In each section we review the field {\\it and} present new results obtained since the last Protostars and Planets meeting.  Where appropriate we also address how new results speak to issues of feedback on star and planet formation. We also attempt to point to ways in which new observing platforms such as ALMA can be expected to influence the field in the near future.  We note that we will focus on outflows from nearby low-mass stars ($<$ 500 pc), which offer the best resolution into the relevant processes. An excellent  review of outflows from high-mass sources was presented in \\citet{2007prpl.conf..245A}, which showed that some (but not all) appear as scaled-up versions of the low-mass case \\citep[see also e.g.][]{2013A&A...550A..81C,2013ApJ...767...58Z}. ALMA will revolutionize our view of these distant and tightly clustered objects so much that our current understanding will greatly evolve in the next few years. Finally we note that our review includes, for the first time, results on jet dynamics from the emerging field of High Energy Density Laboratory Astrophysics (HEDLA).  These experiments and their theoretical interpretation provide direct insights into the 3-D dynamics of fully magnetized, hypersonic, radiative outflows.      \\begin{figure*}[t] \\epsscale{1.5} \\plotone{Hartigan_Summary-eps-converted-to.pdf} \\caption{\\small A schematic view of jets and outflows across seven orders of magnitude in scale.  Note the presence of the scale bar in each figure as one moves from the physics of launching near the star out to the physics of feedback on cluster and cloud scales. See text for reference to specific processes and classes of objects} \\label{Hartigan} \\end{figure*} ",
        "conclusions": "In this chapter we have attempted to demonstrate that protostellar jets and outflows are not only visually beautiful and important on their own as examples of astrophysical magneto-fluid dynamical processes, but they are also an essential player in the assembly of stars across a remarkable range of size-scales. We find that issues of feedback from jets/outflows back to the star formation process is apparent on three scales: those associated with planet formation, those associated with the natal core, and those associated with clustered star formation.  On scales associated with planet-forming disks, jets can impact planet assembly through disk irradiation/shielding and MHD effects associated with outflow launching.  These processes can alter disk properties in those regions where planets will be forming. Future work should focus on articulating the feedback between jet driving and the disk mechanisms associated with creating planets.  On the scales of the natal cores, jets and outflows seem capable of explaining the low observed 30\\% core-to-star efficiency, through a combination of powerful MHD ejection during the earliest collapse phase, and envelope clearing by wide-angle winds during the later phases.  On scales associated with clusters (or perhaps clouds) multiple jets/outflows can drive turbulence, alter star formation efficiencies and affect the mean stellar mass. It is also possible that outflows may help unbind cluster gas.  Future work should focus on providing simulations with the best parametrizations of jet properties (such as momentum injection history and distributions), characterization of physical mechanisms for feedback, and exploration of feedback observational signatures. The full range of scales over which both collimated and uncollimated outflows from young stars impact the star-formation process must also be articulated.  Observational signatures of such feedback must also explored.  With regard to the physics of jets themselves, new results make it clear that they are collimated magnetically on inner disk scales, and include multiple thermal and chemical components surrounded by a {\\em slower} wide-angle wind. The emerging picture is that of a powerful MHD disk wind, collimating the inner stellar wind and magnetospheric ejections responsible for braking down the star. Detailed analysis and modeling is needed to confirm this picture and articulate the properties of these different components as a function of source age.  We have also shown how High Energy Density Laboratory Astrophysics (HEDLA) experiments have already contributed new and fundemnetal insights into the hydrodynamic and MHD evolution of jets.  We expect HEDLA studies to grow beyond jet research in the future as they hold the promise of touching on many issues relevant to star and planet formation (i.e. cometary globules and hot Jupiters).  Observationally we expect new platforms to hold great promise for jet and outflow studies.  In particular ALMA and NIR IFUs should prove crucial to resolving jet rotation profiles, shocks and chemical stratification in statistically relevant jet samples, and to better understand their interaction with the surrounding envelope. Such data will provide definitive tests of disk wind models.  NIR interferometry of CTTS (eg. with GRAVITY on VLTI) promises to be a powerful test of atomic jet models. Synchrotron studies with eVLA, LOFAR should allow jet magnetic fields to finally come into view.  Finally, long baseline monitoring of the short quasi-periodic knot modulation in jets ($\\sim 3-15$ yrs) should allow to clarify the origin of these features and their link with stellar and disk physics (magnetic cycles, accretion outbursts) and source binarity.   \\textbf"
    },
    "1402/1402.3848_arXiv.txt": {
        "abstract": "It is demonstrated, owing to the nonlinearity of QED, that a static charge placed in a strong magnetic field\\ $B$\\ is a magnetic dipole (besides remaining an electric monopole, as well). Its magnetic moment grows linearly with $B$ as long as the latter remains smaller than the characteristic value of $1.2\\cdot 10^{13}\\unit{G}$ but tends to a constant as $B$ exceeds that value. The force acting on a densely charged object by the dipole magnetic field of a neutron star is estimated. ",
        "introduction": "In the recent papers \\cite{ShaGit2012,CGS}, we started a study of self-interaction of static electric and magnetic fields of moderate strength -- in the vacuum \\cite{CGS} or taken against a very strong background formed by a constant and homogeneous magnetic field \\cite% {ShaGit2012}. By \"moderate\" we mean the fields that are strong enough to make the nonlinearity of QED actual for their self-interaction, but still assumed smaller than the characteristic value $B_{\\text{Sch}}=m^{2}/e,$ where $m$ and $e$ are electron mass and charge, to enable us to exploit the expansion of the nonlinear Maxwell equations in powers of the fields and confine ourselves to the lowest terms in this expansion. On the contrary, the background field is not limited either from below or from above. It was observed that effects of selfinteraction manifest themselves already within the simplest approximation of local action, valid in the infrared region of momenta, i.e. for fields that do not change any essentially at the time or space interval of $m^{-1}.$ Among these effects are nonlinear corrections to the Coulomb field in the vacuum \\cite{CGS}, making the field energy of a point charge finite \\cite{Caio2}, electromagnetic nonlinear renormalization of electric and magnetic dipole moments of mesons and baryons \\cite{CGS}, necessary after the latter are calculated following the theory of strong interaction. In the present paper, we continue the work \\cite{ShaGit2012}, where it was pointed out that the quadratic response of the strong background magnetic field to an applied moderate electric field is purely magnetic, and the nonlinear current induced by an applied electric field \\ was calculated based on the third-rank polarization tensor (a three-photon vertex beyond the mass shell) in the infrared (local) limit. Also, a general expression was given for the magnetic field nonlinearly induced by an electrostatic field. Below we demonstrate that this field is that of a magnetic dipole. The magnetic dipole moment carried by a static spherical charge of finite extension is calculated. We estimate the force acting on this magnetic moment by the inhomogeneous magnetic field of a neutron star and a magnetar. If we admit that in the course of the phase transition of the neutron star to a quark star \\cite{Negreiros,quark stars} highly charged domains of strange matter \\cite{strangelets} may exist, then the observed effect may form a mechanism of charge outflow from the star. ",
        "conclusions": ""
    },
    "1402/1402.6376_arXiv.txt": {
        "abstract": "We study the dust present in 56 Galactic planetary nebulae (PNe) through their iron depletion factors, their C/O abundance ratios (in 51 objects), and the dust features that appear in their infrared spectra (for 33 objects). Our sample objects have deep optical spectra of good quality, and most of them also have ultraviolet observations. We use these observations to derive the iron abundances and the C/O abundance ratios in a homogeneous way for all the objects. We compile detections of infrared dust features from the literature and we analyze the available {\\it Spitzer}/IRS spectra. Most of the PNe have C/O ratios below one and show crystalline silicates in their infrared spectra. The PNe with silicates have C/O~$<1$, with the exception of Cn~1-5. Most of the PNe with dust features related to C-rich environments (SiC or the 30 $\\mu$m feature usually associated to MgS) have C/O~$\\ga0.8$. PAHs are detected over the full range of C/O values, including 6 objects that also show silicates. Iron abundances are low in all the objects, implying that more than 90\\% of their iron atoms are deposited into dust grains. The range of iron depletions in the sample covers about two orders of magnitude, and we find that the highest depletion factors are found in C-rich objects with SiC or the 30 $\\mu$m feature in their infrared spectra, whereas some of the O-rich objects with silicates show the lowest depletion factors. ",
        "introduction": "\\label{sec:intro}  Planetary nebulae (PNe) are suitable sites to study the evolution of dust grains since their progenitors, asymptotic giant branch (AGB) stars, are among the most efficient sources of circumstellar dust in the Galaxy \\citep{Whittet_03}. These stars lose large amounts of material and create circumstellar envelopes with both low temperatures (below 1500 K in the regions of dust formation) and high densities ($\\sim$ 10$^{13}$ cm$^{-3}$), where dust grains are efficiently formed.  In \\citet{DelgadoInglada_09}, we studied iron depletions into dust grains in a sample of 33 Galactic disk PNe. Iron is an important contributor to the mass of dust grains \\citep{Sofia_94}, and its gaseous abundance can be used as a proxy for the amount of depletion of other elements. We found low iron gaseous abundances in all the objects, suggesting that more than 90\\% of their iron atoms are located in grains. We also found that the range of iron depletions covers about two orders of magnitude, which may be a consequence of differences in the formation and evolution of the grains in the different nebulae. We did not find any obvious correlation between iron abundances and parameters related to the nebular age (such as the nebular radius or the surface brightness) or to the dominant chemistry in the nebulae (C-rich or O-rich) that could provide clues on the reasons behind the different depletions. However, the information available for the studied sample was scarce.  On the other hand, the dust features that appear in the infrared spectra of PNe provide information on the composition of their dust grains. In our galaxy, PNe with crystalline silicates, many of them also with PAHs, are the most numerous \\citep{Gutenkunst_08, PC_09, Stanghellini_12}. Besides, a general correlation has been found between infrared features and C/O abundance ratios: Magellanic Clouds PNe with silicates have C/O~$<1$, and those with carbonaceous dust have C/O~$>1$, although for Galactic PNe the situation is less clear \\citep{Casassus_01a,Cohen_05,Stanghellini_07,Waters_98}.  The relative abundances of carbon and oxygen in the envelope of an AGB star depends on the nucleosynthesis processes occurring in the inner regions, which in turn depends on the mass of the progenitor star and on its initial metallicity. In our galaxy, we expect stars to be born with an almost solar C/O ratio (C/O~$\\sim0.5$, \\citealt{AllendePrieto_02}), that may change during the evolution of the star depending on its initial mass. Theoretical models \\citep[see, e.g.,][]{Marigo_03, Karakas_09} predict that the less massive progenitors (M$_{*}$~$\\lesssim1.5$--2 M$_\\odot$) will  remain O-rich during their whole evolution. Stars with masses in the range  $\\sim$ 2--4 M$_\\odot$ may evolve from O-rich to C-rich due to the third  dredge-up process that enriches the surface of the AGB star with carbon. In the most massive progenitors (M$_{*}$ $\\gtrsim$ 4--5 M$_\\odot$), the hot bottom burning process occurs efficiently and, since it converts carbon into nitrogen via the CN cycle, it can prevent the formation of a C-rich star.  The value of C/O in a dust-forming stellar envelope defines whether the dust particles will be C-rich or O-rich. If carbon is more abundant than oxygen, all the oxygen atoms will be trapped in CO molecules and carbon atoms will be available to form C-rich compounds, such as graphite, amorphous carbon, or silicon carbide (SiC). If oxygen is more abundant than carbon, all the carbon atoms will be trapped by CO and the remaining oxygen atoms will form O-rich grains like silicates and oxides.  Oxygen rich nebulae could have iron deposited in metallic iron grains, silicates and oxides, whereas C-rich nebulae are expected to have their iron atoms in the form of metallic grains, Fe$_3$C, FeSi, FeS, and FeS$_2$ \\citep{Whittet_03}. Therefore, iron depletion factors could be different in C-rich and O-rich PNe. The C-rich or O-rich chemistry may be studied from the C/O abundance ratios or from the infrared dust features, and we expect both to be related. In principle, C-rich features such as SiC are expected in PNe with C/O~$>1$ \\citep[but see][]{Bond_10} whereas O-rich features like silicates are expected in PNe with C/O~$<1$. {\\it ISO} and {\\it Spitzer} telescopes allow us to detect and identify several dust species in the spectra of PNe and therefore, to classify PNe as C-rich or O-rich according to the observed features.  Here, we extend the work performed in \\citet{DelgadoInglada_09} by studying a sample of Galactic PNe that nearly doubles the previous one and also by analyzing possible correlations between their iron depletions, C/O abundance ratios, and infrared dust features.  The paper is organized as follows. In Section 2 we describe our sample of Galactic PNe, whereas in Section 3 we present the calculation of electron temperatures and densities for all the objects. Section 4 contains an extensive discussion of the iron abundance and depletion factor calculations. Section 5 details the calculation of C/O abundance ratios from collisionally excited lines (CELs) and recombination lines (RLs), and the comparison between these two estimates. The compilation of infrared dust features observed in the sample is presented in Section 6. In Section 7 we discuss our results and study the correlation between iron abundances, C/O values, and dust features in our sample of PNe. Our conclusions are presented in Section 8. ",
        "conclusions": "We have studied a sample of 56 PNe, belonging to the Galactic bulge, the halo, and the disk. The nebulae have available optical spectra of good quality and we constrain their iron abundances using [\\ion{Fe}{3}] lines and the ionization correction scheme derived by \\citet{Rodriguez_05}. The iron depletion factors we find imply that most of the studied PNe have less than 10\\% of their iron atoms in the gas, the missing iron atoms are presumably condensed into dust grains.  We have calculated the C/O abundance ratios using both optical RLs (for 49 PNe) and ultraviolet and optical CELs (for 39 PNe). The differences between the two estimates reach $\\sim0.7$ dex in some objects. Most of the sample PNe show C/O~$<1$ which, according to theoretical models, indicates that they could either descend from stars with masses below $\\sim1.5$ M$_{\\odot}$ or from stars with masses above 4--5 M$_{\\odot}$. Those PNe with C/O~$>1$ are probably the result of intermediate mass stars.  We have also studied the infrared spectra of the PNe with available data in order to identify the following features: PAHs, SiC, the broad band at 30 $\\mu$m, and amorphous and crystalline silicates. We find PAHs in 12--14 of the 33 PNe with infrared spectra. The presence of these features is not restricted to PNe with C/O~$>1$. Silicates are found in 13--16 PNe; SiC and/or the 30 $\\mu$m feature, which is also associated with C-rich environments, in 7--9 objects. The presence of silicates and SiC or the 30 $\\mu$m feature does not agree in all cases with the values of C/O derived from emission lines, maybe reflecting the uncertainties related to the determination of the C/O abundance ratio. Note also that the C/O value we measure can be affected by the depletion of C or O atoms into dust grains.  The oxygen abundances, $12+\\log(\\mbox{O}/\\mbox{H})$, found for the PNe that have different dust features cover the ranges: 8.26--8.87 (PAHs), 8.26--8.68 (SiC or 30 $\\mu$m feature), 8.02--8.87 (silicates), and 8.45--8.87 (mixed-chemistry objects: silicates and PAHs). Therefore, the sample PNe with silicate dust grains seem to arise from a broad population that includes our halo PN, PNe from the Galactic disk, and objects that might belong to the bulge. Our PNe with SiC and the 30 $\\mu$m feature seem to arise from a more homogeneous population in the Galactic disk, while the mixed-chemistry objects arise from PNe of relatively high metallicities, results that agree with those found by \\citet{Stanghellini_12} using measurements of the dust temperature and the distribution in the Galaxy of their sample objects.  We find a relation between the iron depletions and the type of chemistry in the nebulae, C-rich or O-rich. Both C-rich and O-rich PNe cover a wide range of iron depletions, but the PNe with the highest iron depletions have C-rich dust features (SiC and/or the 30 $\\mu$m band) whereas those PNe with the lowest iron depletions have silicates in their infrared spectra. In accordance with this result, we find a weak but significant anticorrelation between the iron abundances and the C/O abundance ratios derived from emission lines. Some kind of correlation is expected from the different molecules and dust compounds that should form in environments with different values of C/O."
    },
    "1402/1402.6695_arXiv.txt": {
        "abstract": "Using idealized 1-D Eulerian hydrodynamic simulations, we contrast the behavior of isolated supernovae with the superbubbles driven by multiple, collocated supernovae. Continuous energy injection via successive supernovae going off within the hot/dilute bubble maintains a strong termination shock. This strong shock keeps the superbubble over-pressured and drives the outer shock well after it becomes radiative. Isolated supernovae, in contrast, with no further energy injection, become radiative quite early ($\\lesssim 0.1$ Myr, 10s of pc), and stall at scales $\\lesssim 100$ pc. We show that isolated supernovae lose almost all of their mechanical energy by a Myr, but superbubbles can retain up to $\\sim 40\\%$ of the input energy in form of mechanical energy over the lifetime of the star cluster (few 10s of Myr). % These conclusions hold even in the presence of realistic magnetic fields and thermal conduction. We also compare various recipes for implementing supernova feedback in numerical simulations. For various feedback prescriptions we derive the spatial scale below which the energy needs to be deposited for it to couple to the interstellar medium (ISM). We show that a steady thermal wind within the superbubble appears only for a large number ($\\gtrsim 10^4$) of supernovae. For smaller clusters we expect multiple internal shocks instead of a smooth, dense thermalized % wind. ",
        "introduction": "Gravity and dark energy govern the structure in the universe at the largest scales but complex baryonic processes like cooling, heating, self-gravity and  star formation are important at galactic scales (e.g., \\citealt{spr05}). Numerical simulations have made tremendous progress in understanding galaxy formation, starting from pure gravitational N-body simulations to the current models which try to model the aforementioned complex processes. Modeling the gravitationally-interacting dark matter is straightforward in principle, and only limited by the available computing power. But the modeling of baryonic processes is rather involved. In particular, there is no consensus on which baryonic processes are important and how they should be implemented numerically. Given the dynamic range of scales, from large scale structure (10s of Mpc) to an individual star forming cloud ($\\sim$pc), simulations have to resort to unresolved `subgrid' models for star formation and feedback due to star formation (e.g., \\citealt{nav93,ger97,spr03,gue11,hop12} and references therein). While different star formation recipes seem to give similar star formation histories and stellar mass distributions, provided molecular clouds are resolved (e.g., \\citealt{hop11}), simulations are quite sensitive to the various feedback prescriptions (e.g., thermal feedback due to supernovae, momentum injection via dust absorbing/scattering photons produced by massive stars and supernovae) even with high resolution.  Stars form in clusters and super-star-clusters (100s to $10^6$ stars) of various sizes and in different environments, ranging from low density galactic outskirts to dense galactic centers (see \\citealt{pro10} for a review). These clusters are observed to disrupt the dense molecular clouds in which they are born (e.g., see \\citealt{lei89,zha01}). This stellar feedback (due to strong radiation, stellar winds and supernovae) disperses cold gas and suppresses further star formation. Because of the stellar initial mass function (IMF) and the main sequence lifetimes, the energy input is roughly constant per unit time over the life time of stars more massive than $8 \\msun$ ($\\sim 50$ Myr; \\citealt{mcc87}). Therefore, stellar feedback is sometimes  modeled as a constant luminosity driven blast wave (\\citealt{wea77,mac89,ger97}). A superbubble (SB) expands faster than an isolated supernova remnant (SNR) because of continuous energy injection, and suffers smaller cooling losses because most supernovae (SNe) go off in a low density bubble.  The hot bubble breaks out through the gas disk if the outer shock driven by overlapping SNe crosses the scale height with a sufficient Mach number within the starburst lifetime (e.g., see section 2 of \\citealt{roy13}). After breakout the hot, metal-rich stellar ejecta is spread out into the galactic halo via the Rayleigh-Taylor instability. Spreading of metals over large scales is required to explain the high metallicity observed in the intergalactic medium (IGM), far away from the stellar disk (\\citealt{tum11}).  Standard models for feedback through multiple SNe assume that a fraction $\\gtrsim 0.3$ of the total explosion energy is retained in the hot ISM (e.g., \\citealt{strickland09}). This fraction is much larger than the estimates for a single SNR ($\\sim 0.1$) after the radiative phase at $\\lesssim 0.1$ Myr (e.g., \\citealt{cox72, che74,tho98}). Moreover, there are further radiative and adiabatic losses, such that over the timescale of order tens of Myr, the available mechanical energy in the bubble and shell is a negligible ($\\ll 0.1$)  fraction of the initial supernova (SN) energy input. Recently \\citet{nath13} argued that SBs created by multiple SNe in a star cluster are more effective as a feedback mechanism, in comparison to incoherent SNe. Recent simulations of interacting SNe by \\citealt{vas14} show that the fractional energy retained as thermal energy in the hot ISM can be as large as $\\sim 0.1\\hbox{--}0.3$ only if the explosions are spatially and temporally correlated such that the radiative losses are effectively compensated by new explosions. SBs produced by compact star clusters are expected to satisfy this condition. In this paper we assume all SNe to be coincident, a good approximation if the SNR, when it becomes radiative, encompasses the whole cluster. This makes SBs a more effective feedback agent which is responsible for magnificent galactic outflows.   In addition to elucidating differences between isolated SNe and SBs, we compare various methods of injecting  SN energy in numerical simulations of the ISM. Different feedback processes, such as radiation pressure, photoionization, and cosmic rays, are important in explaining outflows in galaxies of different masses and redshifts (e.g., \\citealt{hop12} and references therein). In this paper we only focus on the most important feedback component, namely, thermal feedback due to SNe. We find that the SN thermal energy must be deposited over a sufficiently small volume for it to create a hot bubble  and to have an impact on the surroundings. For a large energy deposition radius the cooling time is shorter than the thermalization timescale and thermal feedback is artificially suppressed. Most early implementations of SN feedback suffered from this problem (see \\citealt{ger97} and references therein).  The formation of a low density bubble is essential for thermal SN feedback to work because the energy of subsequent SNe is deposited in the low density  bubble and is not radiated instantaneously; cooling is only restricted to the outer shock. While this fact has been appreciated (e.g., \\citealt{gne98,jou06}), we  present {\\em quantitative} conditions for the formation of a strong shock and a hot, dilute bubble for different thermal feedback prescriptions. The main culprit responsible for the inefficacy of thermal feedback, in both Eulerian (e.g., \\citealt{tas06,dub08}) and Lagrangian (\\citealt{spr03,sti06}) simulations, is the lack of resolution. In reality, a SN affects its surroundings starting at small ($\\ll$ 1 pc) scales by launching very fast ejecta ($\\sim 10^4$ km s$^{-1}$) into the ISM (e.g., see \\citealt{che77}). Simulations in which SN energy is not deposited at small enough scales have to resort to artificial measures (turning off cooling for several energy injection timescales, depositing energy into artificial `hot' phase that does not cool, etc.) in order for feedback to have any impact at all on the ISM (e.g., \\citealt{ger97,spr03,dub08}). These measures are necessary for including the effects of supernova heating in large-scale simulations.  The continuous injection of mass and energy by SNe deep within the hot SB is expected to launch a steady wind (as first calculated by \\citet{che85};  hereafter CC85). CC85 obtained analytic solutions for a superwind assuming a constant thermal energy and mass input rate with an injection radius. By modeling realistic SNe as fast moving ejecta within SBs we show that a steady CC85 wind is obtained only if a large number of SNe ($\\gtrsim 10^4$) go off within the star cluster. For a smaller number of SNe, an individual SN's  kinetic energy is not thermalized within a small injection radius and there is an unsteady outflow. This should have implications for works that simply assume a CC85 wind within the SB.  Our paper presents analytic results that can be readily used for the numerical implementation of thermal feedback using various recipes. These analytic criteria are verified and extended using idealized numerical simulations. The fundamental difference between isolated supernovae and superbubbles -- that the former is ineffective on galactic scales -- is highlighted.  The paper is organized as follows. In section \\ref{sec:fb_pres} we describe various ways of implementing thermal feedback due to SNe. Section \\ref{sec:num_sim} describes the numerical setup used to study SNe and SBs. In section \\ref{sec:anal} we present different analytic criteria for feedback to work with a range of feedback prescriptions. We also derive the conditions for obtaining a thermalized CC85 wind within a SB. In section \\ref{sec:results} we present 1-D numerical simulations of different feedback recipes with and without cooling, and compare the results with our analytic estimates. % We also briefly discuss the effects of magnetic fields and thermal conduction. In section \\ref{sec:conc} we discuss the implications of our results on galaxy formation, feedback simulations, and on galactic wind observations. ",
        "conclusions": "\\label{sec:conc} We have obtained several important results in this paper, on both numerical implementation of SN/SB feedback and on differences between isolated SNe and SBs. SBs are a result of spatially and temporally correlated SNe. Since most massive stars are expected to be born in star clusters few 10s of pc in extent (e.g., \\citealt{lar99}), pre-SN stellar winds and first SNe are expected to carve out a low density bubble, which by a fraction of Myr encloses the whole star cluster (see Fig. \\ref{fig:d_T}). Therefore, subsequent SNe happen in the low density bubble and we are in the SB regime of coherent SNe (see \\citealt{roy13,vas14}).  Magnificent galactic outflows, such as M82, are powered by multiple super-star-clusters and the problem of understanding coalescing SBs is important. Star clusters more massive than $10^5 \\msun$ (and hence with $>1000$ SNe) are rare (e.g., see \\citealt{pro10}); therefore, the superbubbles in M82 and in our $N_{\\rm OB}=10^5$ SB model should be considered as giant bubbles driven by 100s of overlapping superbubbles  due to individual star clusters. Indeed, 100s of star clusters have been observed in the central few 100 pc of M82 (\\citealt{oco95}). We note that  vertical stratification is important for the acceleration and assimilation of the metal-rich bubble into the halo. In this paper we consider the idealized smaller-scale problem of the behavior of isolated SNRs and multiple coincident SNe within a SB in a uniform ISM. Some of our most important results are:  \\begin{itemize}  \\item Our most realistic kinetic explosion (KE) models (and other models in which SN energy is {\\em overwritten}), in which the SN energy in kinetic form is overwritten in a small volume, give correct results only if the energy is deposited within a small length scale (see section \\ref{sec:over}); otherwise, energy is overwritten without coupling to the ISM.  This is true even without considering any radiative losses.  \\item With cooling, if feedback energy is deposited within a length scale $r_{\\rm ej}$ larger than the critical values mentioned in section \\ref{sec:cool_cond}, such that the input energy is radiated before it is thermalized, a hot bubble is not formed in the widely used luminosity driven and thermal explosion addition (LD, TEa) models. Thus SN fizzles out at early stages due to artificially large cooling losses.  \\item With insufficient resolution and large ISM densities the bubble fizzles out completely in the luminosity driven (LD) and thermal explosion addition (TEa) models (in which energy is {\\em added} to the ISM; see section \\ref{sec:fb_pres}) and cannot have any effect on the ISM (see Fig. \\ref{fig:fizzling}).  As also pointed out previously (e.g., by \\citealt{cre11,dal12}), early galactic-scale SN feedback simulations failed mainly because of this. However, for a realistic SN (as mimicked by our KE model) a bubble is formed and subsequent SNe occuring within the non-radiative bubble power the radiative outer shock. Another, probably more serious, problem faced by numerical simulations is that the SN energy is not typically put in coherently over a small volume in space and within a short interval. Feedback due to SNe in young star clusters is expected to be coherent and much more effective than a similar number of isolated SNe (see Figs. \\ref{fig:energy}, \\ref{fig:rad_loss_time}). For correlated SNe, isolated SN bubbles overlap to form a superbubble which is powered by the non-radiative termination/internal shocks, long after the outer shock becomes radiative. In contrast, an isolated supernova becomes powerless a bit after ($\\sim 1$ Myr) the outer shock becomes radiative. \\item A smooth CC85 wind within the superbubble is possible only if the number of SNe ($N_{\\rm OB}$) over the cluster lifetime is large (i.e., $N_{\\rm OB} \\gtrsim 10^4$). Only in these cases, individual SNe going off inside the superbubble are able to thermalize within the termination shock. This result has implications for modeling the X-ray output, for example, in individual bubbles blown by star clusters and in the inner regions of galactic outflows, since the CC85 wind structure is often assumed where it may not be valid.  \\item Most of the radiative losses come from the unresolved radiative relaxation layer at the outer shock. The fractional radiative losses from the interior region, concentrated at the contact discontinuity between the shocked ISM and the shocked ejecta, varies between $\\sim 0.001-0.3$, with larger losses occurring at later times. While these radiative layers are unresolved even in our highest resolution simulations, the volume integrated radiative losses in them converge even for a modest resolution.  \\item As compared to isolated SNe, superbubbles can retain a larger fraction of the initial energy of explosions as thermal/kinetic energy of the gas. Isolated SNe are mixed with the ISM soon after they become radiative; by few Myr they are incapable of affecting the ISM at all. While most energy is radiative away (close to 100\\%, and not $\\sim 90$\\%, as is often assumed) for isolated SNe over 10 Myr, a SB can retain a fraction  $\\sim 0.35$ (for $n=1$ cm$^{-3}$) as the bubble thermal energy + the shell kinetic energy. This fraction is only weakly affected by a higher ISM density and by thermal conduction (see the right panel of Fig. \\ref{fig:rad_loss_time}). Thus, SBs are expected to significantly affect even a dense ISM. Substantial radiative losses can partly explain the smaller observed bubble sizes compared to what is expected by modeling the stellar populations (see \\citealt{oey09} for a summary).  \\item The temperature profile of SBs strongly depend on thermal conduction, whose inclusion can decrease (increase) the temperature (density) and thereby enhance the X-ray luminosity. Thermal conduction (and other sources of mass loading of the hot bubble, such as turbulent mixing) plays an important role in explaining the X-ray emission from galactic superbubbles because very little gas is expected to be in the X-ray emitting regime ($10^6-10^8$ K) in its absence (see the right panel of Fig. \\ref{fig:dT_B_cond}). \\end{itemize}  Our simple one-dimensional simulations show that isolated supernova remnants, owing to large radiative losses, are much weaker feedback agents compared to superbubbles driven by coherently overlapping supernovae. However, detailed three-dimensional calculations, particularly with a realistic distribution of stars in a cluster, and magnetic fields and thermal conduction, are required in order to make quantitative comparisons with observations. This will be done in future."
    },
    "1402/1402.1760_arXiv.txt": {
        "abstract": "We review the appropriateness of using SNIa observations to detect potential signatures of anisotropic expansion in the Universe. We focus on Union2 and SNLS3 SNIa datasets and use the hemispherical comparison method to detect possible anisotropic features. Unlike some previous works where non-diagonal elements of the covariance matrix were neglected,  we use the full covariance matrix of the SNIa data, thus obtaining more realistic and not underestimated errors. As a matter of fact, the significance of previously claimed detections of a preferred direction in the Union2 dataset completely disappears once we include the effects of using the full covariance matrix. Moreover, we also find that such a a preferred direction is aligned with the orthogonal direction of the SDSS observational plane and this suggests a clear indication that the SDSS subsample of the Union2 dataset introduces a significant biased, making the detected preferred direction  unphysical. We thus find that current SNIa surveys are inappropriate to test anisotropic features due to their highly non-homogeneous angular distribution in the sky. In addition, after removal of the highest inhomogeneous sub-samples, the number of SNIa is too low. Finally, we take advantage of the particular distribution of SNLS SNIa sub-sample in the SNLS3 data set, in which the observations were taken along 4 different directions. We fit each direction independently and find consistent results at the 1$\\sigma$ level. Although the likelihoods peak at relatively different values of $\\Omega_m$, the low number of data along each direction gives rise to large errors so that the likelihoods are sufficiently broad as to overlap within 1$\\sigma$. ",
        "introduction": "Recent results from the \\textit{Planck} mission \\citep{Planck} have strengthened the status of the $\\Lambda$CDM model \\citep{LCDM} as the standard cosmological model driving the dynamics of our Universe. This evidence should not be lightly assumed though, as the $\\Lambda$CDM model implicitly requires the validity of a large number of assumptions: the Cosmological Principle \\citep{CosmoPrinciple}, i.e., that our Universe is homogeneous and isotropic on sufficiently large scales; the full validity of General Relativity all the way up to the horizon scale; the existence of unknown dark matter to grow structures via gravitational collapse; the presence of an unnaturally small cosmological constant, which drives the present acceleration of our Universe; and a nearly scale invariant gaussian primordial spectrum of perturbations generated during an inflationary epoch in the early universe. Each of them might be individually questioned \\citep{AgainstLCDM}.  Among them, we will focus here on he validity of the Cosmological Principle and analyse the possibility of detecting a certain amount of anisotropy by using SNIa observations.  Theoretically motivated models giving rise to a violation of the Cosmological Principle by inducing a late-time anisotropic expansion of the universe have been extensively considered in the literature. In \\cite{Campanelli:2006vb}, it was argued that the presence of large scale (homogeneous) magnetic fields \\cite{Campanelli:2006vb} could induce a certain level of eccentricity in the universe expansion that might even solve the low CMB quadrupole problem. The same kind of anisotropic expansion can be achieved by assuming that the dark energy large scale rest frame might differ from that of radiation and/or matter, giving rise to a cosmology with moving fluids. Within this scenario, the dipole acquires a cosmological contribution \\cite{MovingDipole}, the CMB quadrupole is also modified because the relative motion introduces a certain level of eccentricity \\cite{Beltran Jimenez:2007ai} and large scale flows of matter are generated \\cite{Jimenez:2008vs}. The possibility of having different large-scale frames for dark matter and dark energy has also been considered more recently in \\cite{Harko:2013wsa}. Another  scenario with anisotropic expansion was proposed in \\cite{Rodrigues:2007ny}, where the effects of having an anisotropic cosmological constant was studied. More generally,  models with anisotropic dark energy \\cite{Koivisto:2007bp,Battye:2006mb,Campanelli:2011uc,Linder} or anisotropic curvature \\cite{Koivisto:2010dr} have also been considered.  Most of the models  leading to anisotropic expansion discussed in the previous paragraph make the universe metric be of the Bianchi I type, that is, the universe expands at different rates along the 3 spatial directions. Some of those models are characterized by the presence of a privileged direction, which is identified with the vector characterizing the model, direction of the magnetic field, direction of the relative motion between different species. In these cases, the metric is restricted to be of the axisymmetric Bianchi I type in which the expansion rate along the privileged axis differs from the expansion rate along the orthogonal directions. Of course, the amount of anisotropy that these models can generate while being compatible with the highly isotropic CMB is tightly constrained. The main effect comes from the Integrated Sachs-Wolf effect, which is an accumulative effect from the last scattering surface until today. Thus, for models with a non-dynamical evolution of the eccentricity like models with anisotropic equation of state, such anisotropy is essentially constrained to be less than $\\lsim 10^{-4}$ \\cite{Appleby:2009za}. However, if the source of anisotropy is dynamical or there are compensating effects, the constraints are less stringent. In fact, having a period of anisotropic expansion at low redshift would mainly affect the CMB quadrupole, which is affected by a large cosmic variance so that it has less constraining power than the higher multipoles.  Observations of SNIa have been proposed and used as probes of large scale anisotropies. One of the first attempts to constrain the isotropy of the universe by resorting to SNIa measurements was made in \\cite{Kolatt:2000yg}.  In \\cite{Bonvin:2006en} the luminosity-distance dipole from the SNIa distribution was analysed and shown how to use it to obtain measurements of the Hubble expansion rate $H(z)$. Higher multipoles  and additional contributions to the luminosity distance were subsequently computed \\cite{dLanisotropies}. SNIa measurements have also been used to study the isotropy of the Hubble diagram \\cite{Schwarz:2007wf}, \\jose{local bulk flows \\cite{bulkflowsSN}}, the matter distribution \\cite{Antoniou:2010gw} or \\jose{the potential anisotropy of the deceleration parameter \\cite{Cai:2011xs}}.  The possibility of constraining dark energy fluctuations by means of the luminosity distance was explored in \\cite{Blomqvist:2010ky}. \\jose{In \\cite{Mariano:2012wx}, a dipole-like distribution for dark energy was analyzed along with its possible correlation with the fine structure dipole.  Also a dipole-like distribution was considered in \\cite{Cai:2013lja},  but applied to the luminosity distance with respect to the $\\Lambda$CDM case}.  \\jose{A fully Bayesian tool to search for systematic contamination in SnIa data was developed in \\cite{Amendola:2012wc} and further extended, including searches for anisotropic signals, in \\cite{Heneka:2013hka}. }  In the present work we will revisit the constraints that can be obtained on the anisotropy of the universe from SNIa observations. We shall mainly follow the same approach as in \\cite{Antoniou:2010gw} where the authors used a certain version of the hemispherical comparison method to test the isotropy of the matter distribution at the background (homogeneous) level.  We intend to refine their approach by including some influential subtleties concerning the SNIa observations. In particular, we show that including the full covariance matrix with the corresponding increase in the errors makes the significance of previously claimed preferred direction detections disappear. Moreover, we will show that such potential (although not statistically significant) preferred direction happens to be be aligned with the orthogonal direction to the plane of observation of the SDSS subsamples, being an indication of its biased origin due to the particular observation strategy (i.e., of its highly clustered angular distribution along only four directions in the sky) of such a subsample.  The paper is organized as follows:  in \\S.~\\ref{sec:data} we will describe the SNIa samples we have chosen for our analysis and their main properties; in \\S.~\\ref{sec:hemisphere} we will describe the original method we use to test the presence of anisotropic expansion and the main novelties we introduce; in \\S.~\\ref{sec:analysis_pre}  we apply our method to a set of simulated data with a known anisotropic distribution in order to test its validity; in \\S.\\ref{sec:analysis_res} we finally show the results obtained from our analysis and in \\S.\\ref{sec:Conclusions} we conclude by discussing our results. ",
        "conclusions": "\\label{sec:Conclusions} In this work we have presented some new updated results when searching for an anisotropy signal using SNIa measurements. We have used the hemispherical comparison method consisting on fitting opposite hemispheres with independent $\\Lambda$CDM models with the aim of finding an anisotropic distribution of $\\Omega_m$. We have used the Union2 compilation throughout our analysis. Unlike some previous studies, we have considered the full covariance matrix for the SNIa errors and found that it plays a very important role. In fact, considering the full matrix introduces larger errors than those obtained when using only the diagonal errors. Thus, previously reported detections of anisotropy in SNIa observations go away completely  when introducing the existing correlated errors. Moreover, we have used MCMC method, which implies a full detailed angular mapping of the $\\chi^2$ function, allowing to discard possible false minima that could affect the grid or random orientations methods. Our result is in agreement with those in \\cite{Heneka:2013hka}, where no significant indication of hemispherical anisotropy was found. Moreover, we have shown that the particular alignment of the SDSS SNIa seems to introduce a strong bias when searching for anisotropic features. On the other hand, we have used the SNLS3 data and have taken advantage of its very special distribution with the SNIa lying along four different directions. We have then fitted each direction to an independent $\\Lambda$CDM model and have obtained the corresponding matter density parameters. Although the four directions yield likelihoods overlapping at the $1\\sigma$ level, we have found that two of the directions give likelihoods that peak at quite different values of $\\Omega_m$. However, the low number of SNIa in each direction does not allow to draw any significant conclusion.  \\textbf{Acknowledgements} J.B.J. is supported by the Wallonia-Brussels Federation grant ARC No.~11/15-040 and also thanks support from the Spanish MICINN Consolider-Ingenio 2010 Programme under grant MultiDark CSD2009-00064 and project number FIS2011-23000. J.B.J. also wishes to acknowledges the Department of Theoretical physics of the University of the Basque Country for their warm hospitality.  V.S. and R.L. are supported by the Spanish Ministry of Economy and Competitiveness through research projects FIS2010-15492 and Consolider EPI CSD2010-00064, and also by the Basque Government through research project GIC12/66, and by the University of the Basque Country UPV/EHU under program UFI 11/55."
    },
    "1402/1402.5249_arXiv.txt": {
        "abstract": "In a companion paper we have shown how the equations describing gas and dust as two fluids coupled by a drag term can be re-formulated to describe the system as a single fluid mixture. Here we present a numerical implementation of the one-fluid dusty gas algorithm using Smoothed Particle Hydrodynamics (SPH).  The algorithm preserves the conservation properties of the SPH formalism. In particular, the total gas and dust mass, momentum, angular momentum and energy are all exactly conserved. Shock viscosity and conductivity terms are generalised to handle the two-phase mixture accordingly. The algorithm is benchmarked against a comprehensive suit of problems: \\textsc{dustybox}, \\textsc{dustywave}, \\textsc{dustyshock} and \\textsc{dustyoscill}, each of them addressing different properties of the method. We compare the performance of the one-fluid algorithm to the standard two-fluid approach.  The one-fluid algorithm is found to solve both of the fundamental limitations of the two-fluid algorithm: it is no longer possible to concentrate dust below the resolution of the gas (they have the same resolution by definition), and the spatial resolution criterion $h < c_{\\rm s}t_{\\rm s}$, required in two-fluid codes to avoid over-damping of kinetic energy, is unnecessary. Implicit time stepping is straightforward. As a result, the algorithm is up to ten billion times more efficient for 3D simulations of small grains. Additional benefits include the use of half as many particles, a single kernel and fewer SPH interpolations. The only limitation is that it does not capture multi-streaming of dust in the limit of zero coupling, suggesting that in this case a hybrid approach may be required. ",
        "introduction": "\\label{sec:intro}  Gas and dust are the two primary constituents of cold astrophysical systems. Observations of such systems mainly probe the dust phase, with gas properties deduced from the dust distribution assuming a constant $\\simeq 10^{-2}$ dust-to-gas ratio, typical of the interstellar medium. However, the evolution of the gas and dust phases may be quite different, particularly during planet formation. Once the homogeneous mixture assumption breaks down, numerical simulations of the full gas and dust mixture become the only reliable way to predict the properties of the dust phase. Indeed, with the exception of a few problems, the evolution equations for the gas and dust mixture are not tractable analytically. The dust phase is usually treated as a continuous pressureless fluid, an approximation which captures the physics of the system (especially for small grains) but enormously simplifies its description \\citep{Garaud2004}. In particular, the fluid equations are easier to handle compared to integrating the motion of individual grains.  Several numerical methods, both grid-based or Lagrangian, exist to discretise the equations of evolution of a gas and dust mixture. In this paper, we focus on the Smoothed Particle Hydrodynamics (SPH) approach, a fully conservative Lagrangian numerical method (e.g. \\citealt{Monaghan2012,price12}). SPH is well suited to problems with large density contrasts, free boundaries and complex geometries.  Moreover, its Lagrangian nature couples perfectly with the Lagrangian nature of the dust dynamics, making SPH particularly useful for dust-gas mixtures in astrophysics.  \\citet{Monaghan1995} performed the foundational work on a two-fluid SPH gas and dust algorithm, where drag terms were computed by projecting the differential velocity between the particles onto their line of sight to preserve the exact conservation of angular momentum. \\citet{Monaghandust97} improved the method with an implicit timestepping method designed to handle strong drag regimes. This SPH algorithm has been used mostly in the context of planet formation, where grains differentiate strongly from the gas phase \\citep{Maddison2003,Rice04,BF2005,Pinte2007,Maddison2007,Fouchet2007,Laibe2008,Fouchet2010}.  In \\citet{LP12a,LP12b} we highlighted and addressed a number of issues with the existing two fluid SPH algorithm. These included the generalised SPH density estimate in multi-fluid systems, the consistent treatment of variable smoothing length terms and finite particle size, timestep stability, implicit integration, the treatment of non-linear drag regimes, thermal coupling terms and the choice of kernel and smoothing length used in the drag operator. A rigorous mathematical analysis of the algorithm was performed and it was validated against a series of tests, including the analytic solutions recently derived in \\citet{LP11}. The algorithm has been used to predict the filtering and concentration of solids by an existing planet embedded in a protoplanetary disc \\citep{Ayliffe2012}.  Although this improved two-fluid algorithm was found to perform satisfactorily in a number of astrophysical contexts, some important limitations remained \\citep{LP12a,LP12b}. Some of these related to the specific formulation in SPH --- for example, modelling the two fluids as distinct sets of particles meant that interpolation between the two fluids is required. This made it difficult to derive a fast and accurate implicit integration, necessary when the drag is strong, while preserving the exact conservation of angular momentum that is one of the main strengths of the SPH approach \\citep{LP12b}.  More fundamentally, the two-fluid algorithm was found to have two major limitations not specific to SPH \\citep{LP12a}. The first limitation was that both infinite spatial and infinite temporal resolution are required in the rather obvious limit of infinite drag ($\\ts \\to 0$, where $\\ts$ is the stopping time), where the fluids become perfectly coupled. The timestep restriction $\\Delta t < \\ts$ was already well known and can be handled using implicit timestepping methods, as described in \\citep{LP12b}. However, we also found a \\emph{spatial} resolution criterion $\\Delta x \\lesssim \\cs \\ts$, where $\\cs$ is the gas sound speed, in order to prevent artificial separation and thus overdamping of the two-fluid mixture. Since $\\ts$ is very small for strong drag (small grains), this implies a prohibitive computational cost, especially for three dimensional simulations the additional expense may be a factor of several thousand or even millions for micron to millimetre-size grains in protoplanetary discs).  The second limitation was that the total resolution of the simulation is fixed by the more poorly resolved phase. As a result, the structures produced when the dust over concentrates with respect to the gas were found to be strongly resolution-dependent and hence artificial. This prevents the study of systems in which dust is expected to concentrate by orders of magnitude, such as planet formation.  These limitations, which we see as fundamental, have been ignored in almost all gas and dust simulations to date. They are intrinsic to the two-fluid description of the gas and dust mixture. We therefore radically changed our line of attack on the problem. In a companion paper (\\citealt{LP14a}, hereafter \\citetalias{LP14a}) we have proposed a formulation where the gas and dust are described as a \\textit{single fluid made of two different phases}, referred to as \\textit{the mixture}, instead of two fluids interacting with each other. The equations describing the evolution of the mixture were derived. The mixture's mass is the sum of the masses of the two phases and is advected at the barycentric velocity of the two phases. The differential velocity between the two phases and the dust-to-gas ratio, or equivalently the dust fraction, appear naturally as \\textit{internal} properties of the mixture. Treating the dust and gas system as a single fluid naturally solves the two fundamental difficulties described above: In the strong drag limit the equations reduce to the usual fluid equations with a modified sound speed. Secondly. the one-fluid approach implies one resolution length for the system as a whole, meaning that no resolution criterion is required. The one-fluid equations were also found to provide a better physical insight into the behaviour of the mixture compared to their two-fluid counterparts.  In this paper, we derive and test an SPH algorithm for dust-gas mixtures based on the single fluid description. This algorithm is conservative \\textit{by construction}: it conserves the mass, the momentum, angular momentum and the energy of the mixture exactly, as well as the specific mass of each phase taken individually. Furthermore, each SPH particle carries information regarding both phases, removing the need for additional interpolation and/or complicated timestepping schemes.  In Section~\\ref{sec:onefluid} we summarise the relevant parts of the one-fluid formulation from \\citetalias{LP14a}. In Sect.~\\ref{sec:SPH} we derive the SPH versions of these equations. We also discuss timestepping and the necessary modifications to the SPH shock capturing terms. The algorithm is validated against a suite of tests in Sect.~\\ref{sec:tests}, comparing the results against analytic solutions and to those obtained with our previous two-fluid algorithm. We summarise the advantages of the one-fluid approach in Sect.~\\ref{sec:conclu}, and also point out some limitations of this current version of the algorithm in the limit of weak/zero drag. ",
        "conclusions": "\\label{sec:conclu}  We have developed an SPH one-fluid gas and dust algorithm that is fully conservative. The gas and dust are treated as the two phases of the same fluid, the dust-and-gas mixture. This fluid has the total mass of its two phases and is advected at the barycentric velocity. The dust fraction and the differential velocities between the two phases are internal properties of the mixture. The resulting algorithm has been benchmarked against a suit of problems including \\textsc{dustybox}, \\textsc{dustywave}, \\textsc{dustyshock} and \\textsc{dustyoscill}.  The one-fluid algorithm was found to perform particularly well in the situation for which it was designed --- strong drag --- i.e. when dealing with the dynamics of small dust grains. In this situation it was found to be both more efficient (by a factor of $\\sim 10^{10}$!) and more accurate than the standard approach of two fluids coupled by a drag term. Specific advantages compared to the two-fluid approach were found to be as follows: \\begin{enumerate} \\item The spatial resolution criterion $h < \\cs\\ts$ found by \\citet{LP12a} for two-fluid simulations is no longer required; \\item Implicit time integration can be easily implemented and the resulting scheme is very efficient; \\item It avoids resolution issues when dust over-concentrates with respect to the gas, since there is only one resolution length in the mixture; \\item Interpolation errors caused by interpolation between two separate sets of particles are avoided; \\item A factor of two fewer particles are needed; \\item Only one smoothing length is required; \\item The standard bell-shaped SPH kernel is used throughout; \\item The formalism can be implemented in a hydro code without requiring special architecture for multiple fluid species. \\end{enumerate}  The main limitation to the one-fluid method was that the method breaks when the fluid approximation breaks down, i.e. when the velocity field becomes multi-valued. For the most general case where simulations involve mixtures of both large grains and small grains, this suggests that a hybrid scheme might be the best approach. This would be straightforward to implement, consisting of a one-fluid mixture for the small grains, coupled to a separate large grain population treated via the usual two-fluid approach.  The main conclusion of this paper is that the one-fluid algorithm makes possible accurate numerical treatment of strong drag/small grains by removing the fundamental limitations of the two-fluid approach in this regime. Extensions of this study will involve the treatment of a distribution of grain sizes (at the moment, only a single grain size is considered), as well as applications to astrophysical systems where the role of small grains are dominant."
    },
    "1402/1402.2834_arXiv.txt": {
        "abstract": "{ In this paper, we address the migration of small mass planets in 3D radiative disks. Indeed, migration of small planets is known to be too fast inwards in locally isothermal conditions. However, thermal effects could reverse its direction, potentially saving planets in the inner, optically thick parts of the protoplanetary disc.  This effect has been seen for masses larger than 5 Earth masses, but the minimum mass for this to happen has never been probed numerically, although it is of crucial importance for planet formation scenarios. We have extended the hydro-dynamical code FARGO  to 3D, with thermal diffusion. With this code, we perform simulations of embedded planets   down to 2 Earth masses. For a set of discs parameters for which outward migration has been shown in the range of $[5, 35]$ Earth masses, we find that the transition to inward migration occurs for masses in the range $[3, 5]$ Earth masses. The transition appears to be due to an unexpected phenomenon:  the formation of an asymmetric cold and dense finger of gas driven by circulation and libration streamlines. We recover this phenomenon in 2D simulations where we control the cooling effects of the gas through a simple modeling of the energy equation.} ",
        "introduction": "The cores of giant planets and possibly many super-earths form and migrate whitin a protoplanetary disc. In a locally isothermal disc, planets with masses $m_p \\la 10-15$ Earth masses ($M_{\\earth}$) are expected to undergo type-I migration towards the central star on a timescale proportional to $1/m_p$  (\\citet{Tanaka02}) , which can be shorter than the disc lifetime for $m_p \\ga 1 M_{\\earth} $. If this were true, planetary cores should not survive in the disc,  in evident contrast with the observation that many planetary systems, including our own, have planets that are not even very close to the star. \\par Many efforts have been done to solve this puzzling problem. Different mechanisms have been suggested  to halt inward migration such as magnetic field effects in turbulent discs (\\citet{Baruteau11,Guilet13} and references therein) or density trap (\\citet{Masset02,MassetMorby06,MorbyCrida08}). But, perhaps the most promising mechanism appears when considering the effect of energy transfer within the disc.  Studies have  shown that non-isothermal effects can change the direction of the type-I migration (\\citet{PaardeMelle06,BarMass08,KleyCri08,KBK09,MasCas09,MassetCasoli10,Paaretal10}). \\par Precisely,  the migration in the inner part of a  radiative disc can be directed outward, while it remains directed  inward in the outer disc. This establishes the existence of a critical radius where migration vanishes, towards which  planetary cores  migrate from both the inner and the outer part of the disc. Therefore, the zero  migration location acts as  a planet trap (\\citet{Lyraetal10}). \\par The   torque exerted by the protoplanetary disc has two main contributions:  i) the so-called Lindblad torque due to the spiral arms launched by the planet  in the disc which is not affected by the equation of state \\footnote{Actually, it scales with $\\gamma \\Gamma_0$, with $\\Gamma_0$ given in \\ref{gamma0}.},  and ii) the corotation torque caused by material librating in the horseshoe region. When considering radiative effects it has been shown that the corotation torque contribution can be positive and possibly dominate over the negative Lindblad torque. In a recent paper, \\citet{Paaretal11}, have provided a detailed formula for the torque which was calibrated with a 2D hydrodynamical model for low mass planets ($5 M_{\\earth}$). However, the corotation effect depends on  the disc parameters and on the mass of the planet.  Decreasing the planetary mass, the horseshoe {U-turn} timescale {increases and, if it} becomes longer than the radiative diffusion timescale{, the} disc behaves similarly to the isothermal case and the positive torque contribution fades. \\par The  aim of this paper is to  investigate  the total torque acting on low mass planets kept on fixed orbits and possibly  find the critical mass for which the positive contribution of the corotation torque becomes smaller than the Lindblad torque. To this purpose we have extended  the FARGO code of \\citet{Masset00}, \\citet{Masset00b} to 3 dimensions  introducing an energy equation to provide a realistic modeling of radiative effects as in \\citet{KBK09}. We have validated the 3 dimensional version of FARGO (FARGO with Colatitude Added, FARGOCA hereafter) by reproducing some results published in the literature (see Appendix) . \\par We have computed the total torque on planets with mass ranges in the interval $[2,10]$ Earth masses for a specific set of disc parameters and we found that a transition from positive to negative total torque occurs for a planet with mass between 3 and 5  Earth masses. We remark that for the same disc parameters \\citet{KBK09} found positive total torque for planetary masses in the range $[5,35]$ Earth masses.  When  applying the torque formula provided in  \\citet{Paaretal11} to our disc parameters we  find a transition from outward to inward migration occurring between  2.7 and 3  Earth masses. \\par Despite the good agreement between model and simulations it is important to study the mechanism that causes the transition  since in the case of \\citet{Paaretal11} the results were calibrated with 2D simulations while our work is done with numerical simulations of a full radiative 3D model. For this purpose we have analyzed the  torque suffered by the planet from every radial ring of the disc.   We have found a  new and unexpected feature:  a negative contribution to the torque that is not seen for planetary masses larger than $m_p \\ga 5M_{\\earth}$. This is due to an overdensity of gas behind the planet location just outside of its orbit.   We have recovered the same effect in 2D simulations where the energy equation has a  simple cooling prescription in which the cooling time appears as a parameter. Despite its simplicity, the 2D model allows us  to confirm the result and to study its dependence on the gas cooling timescale. \\par  The paper is organized as follows: the physical modelling is presented in Section 2, in Sect. 3   total torque calculations  on small mass planets are described. We provide in Sect. 4 a detailed analysis of the local torque acting on the planets as well as a mechanism for the origin of the negative and positive local contributions to the torque  that we observe on small mass planets. We validate the proposed mechanism in Sect. 5 using  2D simulations. In Sect. 6 we explain the origin of the asymmetry between negative and positive torque contributions. Conclusions are provided in Section 7.   In the appendix we provide test calculations  on the FARGOCA code. ",
        "conclusions": "Using 3D hydrodynamical simulations including radiative transfer we have computed the torque acting on small mass planets  kept on fixed orbits in the aim of finding  the critical mass for which the positive contribution of the corotation torque becomes smaller than the Lindblad torque. While we were expecting to observe  that the positive corotation torque fades away for decreasing mass of the planet,  we found a  new and unexpected feature:  a negative contribution to the torque that is not seen for planetary masses larger than $m_p \\ga 5M_{\\earth}$. We explain this phenomenon observed in viscous discs \\footnote{we do not refer to the nonviscous case nor to the limit of vanishing viscosity}  as follows. Small mass planets don't have a circum-planetary disc, but their gravitational force still increases the pressure and temperature in their neighborhood.  Looking at density and temperature fields it appears  that  some circulating or even some librating streamlines go through pressure increase as they encounter the planet. If the equation of state is locally isothermal or purely adiabatic, the temperature, density and pressure are identical before and after this pressure peak. However, if radiative transfer is included, the gas is cooler, thus denser, after the encounter with the planet. This leads to the formation of a cold and dense finger ahead of the planet just inside of its orbit, and behind the planet just outside of its orbit, outside of the horseshoe region. This phenomenon, that we have recovered with two dimensional high resolution simulations, explains the unexpected negative contribution observed on the torque just outside the planet orbit in 3D simulations. \\par At our knowledge the cold finger was not observed in previous  numerical computations. We understand that, in order to have a torque contribution related to the cold finger in viscous discs, the following conditions must be satisfied: \\begin{itemize} \\item { planets have to be small enough  for their Hill radius to be significantly larger than their Bondi radius.} \\item the equation of state must be non  isothermal neither adiabatic, with an appropriate cooling time. \\item the planetary potential has to be close to the real one, for example the cubic potential introduced in \\citet{KBK09} suited to 3D simulations. \\end{itemize}  The cold finger turns out not to be symmetric being more pronounced behind the planet with our standard disc's parameters, and it   appears to be responsible for the transition from outward to inward migration as the planet's mass decreases. However, this asymmetry is very sensitive to the properties of the flow around the planet. In particular, we found a positive total torque in a disc with flat surface density and temperature profile. Therefore, for specific disc properties, earth size planets could still migrate outwards, even when the usual thermal contribution to the corotation torque fades. This enlightens the importance of the cold finger, and advocates for future studies of this new torque contribution."
    },
    "1402/1402.3235_arXiv.txt": {
        "abstract": "{Measuring the stellar position angle provides valuable information on binary stellar formation or stellar spin axis evolution. } { We aim to develop a method for determining the absolute stellar position angle using spectro-astrometric analysis of high resolution long-slit spectra. The method has been designed in particular for slowly rotating stars. We investigate its applicability to existing dispersive long-slit spectrographs, identified here by their plate scale, and the size of the resulting stellar sample. } { The stellar rotation induces a tilt in the stellar lines whose angle depends on the stellar position angle and the orientation of the slit. We developed a rotation model to calculate and reproduce the effects of stellar rotation on unreduced high resolution stellar spectra. Then we retrieved the tilt amplitude using a spectro-astrometric extraction of the position of the photocentre of the spectrum. Finally we present two methods for analysing the position spectrum using either direct measurement of the tilt or a cross-correlation analysis. } { For stars with large apparent diameter and using a spectrograph with a small plate scale, we show that it is possible to determine the stellar position angle  directly within $10^\\circ$ with a signal-to-noise ratio of the order of 6. Under less favourable conditions, i.e. larger plate scale or smaller stellar diameter, the cross-correlation method yields comparable results.} { We show that with the currently existing instruments, it is possible to determine the stellar position angle of at least 50 stars precisely, mostly K-type giants with apparent diameter down to 5 milliarcseconds. If we consider errors of around $10^\\circ$ still acceptable, we may include stars with apparent diameter down to 2 mas in the sample that then comprises also some main sequence stars.} ",
        "introduction": "\\label{sec:Intro}  Angular momentum vectors are usually assumed to be randomly oriented during the star formation process leading to a uniform distribution of stellar spin axis orientations. However, there might be a physical process -- for instance magnetic fields (\\cite{Hennebelle})-- that would lead to a preferred orientation on the scale of a star forming region despite chaotic stellar formation (\\cite{Bate}). As a result, measuring the absolute orientation of the stellar spin axis would provide valuable information for constraining the models for stellar formation and stellar evolution, such as spin axis distribution inside a stellar cluster, spin axes alignment of binaries, and evolution of the stellar spin axis orientation with the age of the star.   The spin axis orientation is defined by two angles: inclination $i$, which is the angle formed between the spin axis and the line of sight with the observer, and the position angle, which is the projection of the spin axis on the sky and measured from north to east. In this paper we concentrate on determining the position angle.   For eclipsing objects, such as binary stars or a star with an exo\\-planet, measuring the Rossiter-McLaughlin effect, i.e. the dimming of light as a function of wavelength, provides additional information on the system geometry. During the eclipse,  a portion of the rotating star is blocked by its companion, causing a weakening in the corresponding Doppler component of the stellar absorption lines. The shape of this spectral distortion depends on the angle between the stellar spin axes, in the case of binaries, or stellar and orbital spin axes, in the case of exoplanets.  This effect is currently exploited in many cases of transiting exoplanets to deliver the spin-orbit alignment. However, the associated error is quite important. To date, the mean error to spin-orbit measurements is close to $10^\\circ$. There are no constraints available for the orientation of the absolute stellar spin axis.   The technical improvements in long baseline interferometry during the past decade have opened the possibility of resolving and imaging stellar surfaces quite accurately. Several teams have observed the absolute spin axis orientation of a dozen early-type stars: Altair (\\citet{Monnier}), Vega (\\citet{Peterson}), and Achenar (\\citet{deSouza}). The spin axis orientation is actually a by-product of determining the gravitational limb darkening coefficient for these fast rotating stars. Gravitational limb darkening causes a shift of the photocentre towards the poles. The amplitude of the  shift is highly dependent on the rotational velocity of the star. For slowly rotating objects, which are mainly late-type stars, the method soon becomes unreliable because the shift is no longer measurable.  \\citet{Petrov} devised an alternative interfero\\-metric method of determining the position angle on large,  slowly rotating stars using differential interferometry. The method combines speckle interferometry and spectroscopy. They show that using a high-to-moderate-resolution spectrograph, coupled with an interferometer, the position angle could be measured within $10^\\circ$ precision. The method was tested successfully on Aldebaran (\\citet{Lagarde}). However, owing to the short speckle integration time, its use remained limited to the brightest stars (Lagarde, private communication).  There is currently a lack of observational methods for determining stellar spin orientation for slowly rotating stars. We investigate the potential of spectro-astrometry to measure the position angle of these stars, with focus on late-type giants. We aim here to define the optimal conditions, instrumental and observational, in order to accurately measure, within $10^\\circ$ errors, the position angle of a large sample of stars and the required signal-to-noise ratio (S/N).    Spectro-astrometry concentrates on the spatial information contained in the spectrum by measuring the wavelength dependent position of the photocentre of an object. It is a powerful tool that allows us to reach sub-diffraction resolution using standard directly-fed spectrographs from otherwise unresolved sources. In addition, it is an instrumentally easy method, since it only requires a stable spectrograph and a detector. However, it also requires a special observing practice, since different slit orientations are required for the analysis. Nonetheless it can be applied to a wide variety of science cases. It has been used successfully in the study of unresolved binaries (\\citet{Bailey, Baines}), of outflows from young objects down to the AU scale (\\cite{Takami}), and of jets from brown dwarfs (\\citet{Whelan}). Furthermore, \\citet{Voigt} applied the method for mapping stellar spots on red giant surfaces using \\textsc{Crires}\\footnote{Cryogenic high Resolution Infrared \\'Echelle Spectrograph}.   In a long-slit spectrograph, the spatial information is conserved from the slit to the detector. Indeed the spatial scale on the detector can be directly derived from the optics of the spectrograph. The position of the photocentre is determined either by fitting a 1D Gaussian over the order or by a weighted arithmetic mean for each wavelength. As a result, even if the source is unresolved, as long as it presents an asymmetric spectral energy distribution, the position of the photocentre will diverge from the continuum at certain wavelengths. Furthermore, it can be measured to sub-pixel precision. The only limitation is set by the pixels inhomogeneities of the detector and the signal-to-noise ratio. The precision reached in the position of the photocentre is approximated by \\begin{math} \\sigma = 0.5 \\times \\text{FWHM}\\times \\text{S/N}^{-1} \\end{math}.  For instance, with an average seeing of 1\", assuming a detector with an one-to-one conversion from photon to electron, a single exposure can achieve up to 40 000 photons counts in the spectrum. Then using the previous formula, the precision in the position is estimated close to 2.5 milliarcseconds (mas). This scale is comparable with the apparent diameter of our closest or largest neighbour stars.   \\vspace{1em} This paper is organised as follows. In Sect \\ref{sec:2ddoppler}, we define a 2D model for  stellar rotation adapted for use in the spectro-astrometric analysis. In Sect 3, we describe the simulation set-up for the construction of the 2D spectrum and the extraction of the spectro-astrometric signal. Section 4 presents the results of our simulations, which are discussed in Sect 5. ",
        "conclusions": "\\label{sec:conclu}  We have shown that stellar rotation induces a sub-pixel tilt in the stellar spectral lines whose inclination depends on the absolute stellar position angle. By using a spectro-astrometric analysis of high resolution spectra, we proved that it is possible to determine the absolute stellar position angle with $10^\\circ$ accuracy for a moderate sample of stars. Our simulations demonstrated two possible ways to retrieve the spectro-astrometric signal, i.e. direct measurement of the photocentre displacement or cross-correlation analysis. The appropriate method depends on the signal-to-noise of the spectra, the observing set-up, i.e. the spectrograph's plate scale, the target, i.e. the apparent stellar diameter, and the slit width. However, we have shown that there is a optimal combination between plate scale and stellar radius that optimizes the detection rates for a constant raw signal-to-noise ratio.  According to our results, with an instrument like CRIRES, which provides a plate scale of 0.086\"/pix when coupled with an adaptive optic, it is possible to determine the position angle of stars with apparent diameter down to 2 mas with errors that are lower than $10^\\circ$. The observations should be taken for at least four different slit angles, and can be obtained either sequentially with seeing effects, or simultaneously using additional instrumentation. This represents a sample of about 100 targets, from A-type main sequence to various K-type giants, including a handful of known spectroscopic binaries. Probing the stellar position angle for these targets would provide a first step in answering two of the original questions: can the stellar spin axis change with the stellar evolution, and how well are the spin axes of binaries aligned."
    },
    "1402/1402.4309_arXiv.txt": {
        "abstract": "Origin of the TeV gamma ray emission from MGRO J2019+37 discovered by the Milagro experiment is investigated within the pulsar wind nebula (PWN) scenario using multiwavelength information on sources suggested to be associated with this object. We find that the synchrotron self-Compton (SSC) mechanism of origin of the observed TeV gamma rays within the PWN scenario is severely constrained by the upper limit on the radio flux from the region around MGRO J2019+37 given by the Giant Metrewave Radio Telescope (GMRT) as well as by the x-ray flux upper limit from SWIFT/XRT. Specifically, for the SSC mechanism to explain the observed TeV flux from \\mgro\\ without violating the GMRT and/or Swift/XRT flux upper limits in the radio and x-ray regions, respectively, the emission region must be extremely compact with the characteristic size of the emission region restricted to $\\lsim{\\mathcal O}(10^{-4}\\pc)$ for an assumed distance of $\\sim$ few kpc to the source. This is at least four orders of magnitude less than the characteristic size of the emission region typically invoked in explaining the TeV emission through the SSC mechanism within the PWN scenario. On the other hand, inverse Compton (IC) scattering of the nebular high energy electrons on the cosmic microwave background (CMB) photons can, for reasonable ranges of values of various parameters, explain the observed TeV flux without violating the GMRT and/or SWIFT/XRT flux bounds. ",
        "introduction": "The Cygnus region of the Galaxy hosts a number of extended, unidentified sources of TeV $\\gamma$-ray emission, the most prominent of which is \\mgro\\  discovered by the Milagro experiment~\\cite{2007ApJ...658L..33A,2007ApJ...664L..91A}. Detailed analysis~\\cite{2012ApJ...753..159A} of the observational data on this object collected during the period 2005--2008 gives a detection of this source with a statistical significance in excess of $12\\sigma$ between 1 and 100 TeV. The measured flux~\\cite{2012ApJ...753..159A} from this source is $7^{+5}_{-2}\\times10^{-10}\\s^{-1}\\m^{-2}\\tev^{-1}$ (with a $\\sim30\\%$ systematic uncertainty) at 10 TeV with a spectrum that is best described by a power-law with a spectral index of $2.0^{+0.5}_{-1.0}$ (with a systematic uncertainty of $\\sim 0.1$) and an exponential cutoff at an energy $E_c=29^{+50}_{-16}\\tev$.  \\paras Although no confirmed counterparts of the TeV source \\mgro\\ at lower energies are known, several possible associations with other observed sources have been suggested. The emission from \\mgro\\ may be due to either a single extended source or several unresolved sources. The EGRET sources 3EG J2021+3716 and 3EG J2016+3657 are positionally close to \\mgro, and thus could be the GeV counterparts of \\mgro\\ if it is a multiple source. At the same time, the EGRET source 3EG J2021+3716 is suggested to be associated with the radio and GeV pulsar PSR J2021+3651 (and its associated pulsar wind nebula PWN G75.2+0.1 \\cite{2002ApJ...577L..19R,2004ApJ...612..389H}) observed at GeV energies by AGILE (AGL J2020.5+3653)~\\cite{2008ApJ...688L..33H} as well as FERMI (2FGL J2021.0+3651)~\\cite{2009ApJS..183...46A}. A SWIFT/XRT observation~\\cite{2007ATel.1097....1L} was also done within the positional uncertainty region of \\mgro\\ reported in \\cite{2007ApJ...658L..33A,2007ApJ...664L..91A} and three x-ray sources were reported in the region with a total x-ray flux corresponding to $\\nu F_\\nu \\sim 8.1\\times10^{-14}\\tev\\,\\cm^{-2}\\,\\s^{-1}$ in the 2--10 keV energy region, which can, therefore, be taken as an upper limit on the possible x-ray flux from any x-ray counterpart of \\mgro\\  in this energy region. In addition, a wide-field deep radio survey of the \\mgro\\ region at 610 MHz was made by the Giant Metrewave Radio Telescope (GMRT)~\\cite{2009A&A...507..241P}, yielding no detectable radio source, thus giving a conservative upper limit of $\\sim$ 1.0 mJy on the radio flux from any point-like radio counterpart of \\mgro. Recently, a deep very high energy (VHE) observation by the VERITAS experiment \\cite{0004-637X-788-1-78} has resolved the VHE emission from the region of MGRO J2019+37 into two VHE sources. One of these, namely, VER J2019+378, which is also positionally coincident with the energetic GeV pulsar PSR J2021+3651 and its PWN, has been suggested to be associated with the bulk of the TeV emission from \\mgro.  \\paras In this paper, we study the implications of a scenario in which the observed TeV \\gray emission from \\mgro\\ arises from a Pulsar Wind Nebula (PWN) type source. Pulsar Wind Nebulae (PWNe) (see, e.g., \\cite{2006ARA&A..44...17G} for a review), a well-known example of which is the Crab Nebula (see \\cite{2008ARA&A..46..127H} for a review), are known to be sources of very high energy gamma rays extending to TeV energies (see, e.g., \\cite{Aharonian_book_2004,2012ASPC..466..167K} for reviews). The TeV photons are thought to be emitted mainly through (a) inverse Compton (IC) interaction of high energy electrons with the low energy synchrotron photons emitted by the electrons themselves in the ambient magnetic field in the nebula --- the so-called synchrotron self-Compton (SSC) mechanism, and/or (b) IC interaction with the photons constituting the cosmic microwave background (CMB) and infrared background --- hereafter referred to as the ``IC-CMB\" mechanism. The high energy electrons themselves are thought to be accelerated in the wind termination shock where the relativistic wind from the pulsar residing within the nebula is stopped by the nebular material. In principle, in addition to electrons high energy protons (and in general heavier nuclei) may be also accelerated ~\\cite{1996MNRAS.278..525A,1997PhRvL..79.2616B,2003A&A...402..827A,2003A&A...405..689B}, which can produce high energy photons through decay of neutral pions produced in inelastic $p$-$p$ collisions. In this paper we shall restrict our attention to the leptonic scenario, i.e., we assume that the TeV emission is due to electrons. We use multiwavelength data and flux upper limits from observations in the region around \\mgro\\ including the radio upper limit given by GMRT~\\cite{2009A&A...507..241P}, x-ray flux upper limit from SWIFT/XRT observations \\cite{2007ATel.1097....1L}, GeV observations by FERMI~\\cite{2009ApJS..183...46A}, EGRET~\\cite{1999ApJS..123...79H} and AGILE~\\cite{2008ApJ...688L..33H} and the TeV data from Milagro~\\cite{2007ApJ...658L..33A,2007ApJ...664L..91A,2012ApJ...753..159A} and VERITAS~\\cite{0004-637X-788-1-78} to study the implications of both the SSC and IC-CMB mechanisms for the macroscopic parameters of the underlying PWN, namely, its energetics (e.g., the total energy contained in the high energy electrons), the characteristic size of the emission region and the ambient magnetic field within the nebula.  \\paras We find that the SSC mechanism of origin of the observed TeV gamma rays from \\mgro\\ is severely constrained by the GMRT upper limit on the radio flux from the region around MGRO J2019+37 as well as by the SWIFT/XRT x-ray flux upper limit. Specifically, for the SSC mechanism to explain the observed TeV flux from \\mgro\\ without violating the GMRT and/or Swift/XRT flux upper limits, the emission region within the PWN must be extremely compact with the characteristic size of the emission region restricted to $\\lsim{\\mathcal O}(10^{-4}\\pc)$ for an assumed distance of $\\sim$ few kpc to the source. This is at least four orders of magnitude less than the characteristic size of the emission region typically invoked in explaining the TeV emission through the SSC mechanism within the PWN scenario. On the other hand, the IC-CMB mechanism can, for reasonable ranges of values of various parameters, explain the observed TeV flux without violating the GMRT and/or SWIFT/XRT flux bounds.  \\paras The reason for the upper limit on the size of the emission region in the SSC scenario is not hard to understand: Ignoring for the moment the details of the energy spectrum of the electrons, let $n_e$ denote the number density of the electrons and $\\rem$ the characteristic radius of the (assumed spherical) emission region in the source. Recall that in the SSC mechanism the TeV photons are produced through IC interaction of the nebular high energy electrons with the synchrotron photons produced by the electrons themselves in the magnetic field in the nebular region. Since the number density of the synchrotron photons scales as $n_e$, the number density of the TeV photons produced by the electrons through the SSC mechanism roughly scales as $n_e^2$. Thus, for a given distance to the source, the emerging total TeV flux from the source scales as $n_e^2\\, \\rem^3$. The requirement of producing the observed TeV flux of \\mgro, therefore, fixes the product $n_e\\,\\rem^{3/2}$. On the other hand, the photon fluxes in the radio and x-ray regions due to synchrotron emission by the electrons scale with the product $n_e\\,\\rem^3$. Therefore, with the product $n_e\\,\\rem^{3/2}$ fixed by the observed TeV flux of \\mgro, an upper limit on the radio flux given by GMRT or the x-ray flux given by SWIFT/XRT directly yields an upper limit on $\\rem$. These arguments are elaborated upon more quantitatively in the following sections within the context of a simple power-law form of the energy spectrum of the electrons within the nebula.  \\paras The rest of the paper is organized as follows: First, in Section~\\ref{sec:Multiwavelength_spectra} we set up the formulae to calculate the multiwavelength photon spectra and compare the resulting theoretically calculated multiwavelength photon spectra with the observed multiwavelength data and constraints pertaining to \\mgro, and discuss the implications --- in particular the constraints on the characteristic size of the emission region --- for the SSC mechanism of production of the observed TeV flux from \\mgro. Finally, we summarize our results and conclude in Section~\\ref{sec:summary_conclusions}. ",
        "conclusions": "\\label{sec:summary_conclusions} To summarize, in this paper we have considered a PWN scenario of the origin of the observed TeV gamma ray emission from \\mgro. We find that, while no lower energy counterparts of this object have yet been identified, the upper limits on possible radio and x-ray emissions from the source provided by the GMRT and SWIFT/XRT, respectively, already imply that the observed TeV emission of \\mgro\\  must originate from a highly compact region with characteristic size of the emission region $\\lsim{\\mathcal O}(10^{-4}\\pc)$ if the observed TeV flux from \\mgro\\ is dominantly produced by the SSC mechanism. Such a compact size of the emission region is difficult to envisage within the usual PWN scenario. On the other hand, IC scattering of the high energy electrons on the CMB photons can  explain the observed TeV flux without violating the radio and x-ray flux upper limits given by GMRT and SWIFT/XRT, respectively, for reasonable ranges of values of various relevant parameters.  \\parab\\paras {\\bf References}"
    },
    "1402/1402.1140_arXiv.txt": {
        "abstract": "The steadily improving sensitivity of pulsar timing arrays (PTAs) suggests that gravitational waves (GWs) from supermassive black hole binary (SMBHB) systems in the nearby universe will be detectable sometime during the next decade. Currently, PTAs assume an equal probability of detection from every sky position, but as evidence grows for a non-isotropic distribution of sources, is there a most likely sky position for a detectable single source of GWs? In this paper, a collection of galactic catalogs is used to calculate various metrics related to the detectability of  a single GW source resolvable above a GW background, assuming that every galaxy has the same probability of containing a SMBHB. Our analyses of these data reveal small probabilities that one of these sources is currently in the PTA band, but as sensitivity is improved regions of consistent probability density are found in predictable locations, specifically around local galaxy clusters. ",
        "introduction": "\\setcounter{footnote}{0} \\setcounter{figure}{0}  Pulsar Timing Arrays (PTAs) are collections of millisecond pulsars \\citep{Backer90} whose joint timings will show correlations that are a specific signature of gravitational waves (GWs) passing between the Earth and the pulsars \\citep{Sazhin78, Detweiler79}. Currently, there are three PTAs working collaboratively to detect GWs in this way \\citep{Hobbs10}. These arrays will soon have the sensitivity to detect single extragalactic sources of GWs, \\citep{Yardley10, Ellis12} which are resolvable above a GW background \\citep{Sesana13CQG}. To increase the efficiency of PTA observations and potentially decrease the time until a detection is made, various groups have considered the optimization of PTA observations \\citep{Lee08, Burt10, Lee12}, but all have assumed an equal probability of detection across the sky. While all directions are equally likely to contain a GW source, there should exist more probable locations for the brightest GW source. And while the probability is quite small for there to exist a source that stands out above the background, the analysis in this paper identifies the potential locations of that source using what is currently known about the distribution of galaxies in the local universe, and allows others to enhance discussions on the optimization of a PTA \\citep{Anella13}.  Supermassive black hole binary (SMBHB) systems with periods of months to years are thought to be the most important source of gravitational waves \\citep{Jaffe03}. Binaries like these form when galaxies containing nuclear black holes merge \\citep{Begelman80, Volonteri03}. Corresponsingly, the number of such binaries should be greatest where galaxy mergers are more frequent; i.e., in galaxy clusters. All things being equal, the nearest clusters will play host to the brightest sources. This sugests that, as PTA sensitivities increase, we look toward the largest or richest nearby clusters (Virgo, Fornax, Norma, Perseus and Coma) as the most likely location of the earliest detectable SMBHB source. Here we make this expectation quantitative.  Below in \\S 2, we use a compilation of several galaxy surveys to identify the mass, distance, and location of all galaxies within $150~$Mpc.  From this data we estimate the SMBHB mass and the lifetime of GW emission in a detectable PTA band for each galaxy in \\S 3, and together with the distance, we identify the probability of the existence of a detectable source in a given direction. \\S 4 contains a full explanation of the results with sky position maps. We investigate the probability of detection given an increasing PTA sensitivity in \\S 5 and a summary of our findings is found in \\S 6. ",
        "conclusions": "In this paper, data is gathered from a compilation of at least $90\\%$ complete galactic surveys out to 200 Mpc. These galaxies are all assumed to contain a SMBHB with equal mass black holes, and the total central mass is calculated using B-band luminosity as a surrogate for mass. This data set is then cut down to only detectable sources (i. e. sources with a total central mass larger than $10^7$ and with in a distance of 150 Mpc). From this data, two metrics are used to estimate the ``hotness\" of the GW sky: $\\langle N_7 \\rangle$, which is proportional to the number density of sources in a given direction; and $\\langle P \\rangle$, which is proportional to the expectation value of the power in GWs emitted from a particular direction. Using $\\langle N_7 \\rangle$, we are able to calculate the specific probability that each pixel contains a currently radiating GW source in the PTA band. While this number is very small for any individual pixel, we find a total probability of $0.023$ that one of the galaxies we considered is detectable. The `brightest' location on the sky has a $2.9 \\times 10^{-4}$ probability of currently containing a single source that stands out about the background in the PTA band. Overall the distribution of single sources potentially detectable by a PTA has a larger number density around local galaxy clusters. While the GW signal strength is dominated by a handful of sources, with the region of sky around the Virgo cluster having a larger number density of these `bright' sources.  Work by \\citet{Anella13} has been done to show that if a new powerful timing pulsar was discovered in the direction of the Virgo cluster, a PTA would have twice the sensitivity to a region about $20^o$ around the discovered pulsar. Given that the sensitivity of a single pulsar to a GW source falls off as $1 + cos \\theta$, where $\\theta$ is the angle between a pulsar and a GW source \\citep{Burt10}, we recommend focusing the search for new pulsars in the vicinity of the Virgo cluster.  As stated earlier in this paper, this is only a broad estimate that assumed an equal probability for all galaxies to contain a binary with equal mass black holes. Future work is planned to incorporate work done by \\citet{Rosado13}, which uses the Sloan Digital Sky Survey and the Millenium simulation data \\citep{Springel05} to search for SMBHBs in the redshift range of $z = 0.01 - 0.7$. Combining these two data sets is the next step towards creating realistic population distributions for single source   GWs detectable by PTAs. While the distribution of detectable single sources will most likely scale with any new estimate, $\\langle P \\rangle$ is more affected by these specific probabilities, and a future paper will address these factors."
    },
    "1402/1402.6884_arXiv.txt": {
        "abstract": "{Using data derived from the H.E.S.S. Phase 1 telescope system and a Ceilometer facility on site, a method of correcting for changing atmospheric quality based on reconstructed shower parameters is presented. The method was applied to data from the active galactic nucleus PKS 2155-304, taken during August and September 2004 when the quality of the atmosphere at  the site was highly variable. Corrected and uncorrected fluxes are shown, and the method is discussed as a first step towards a more complete atmospheric calibration.} ",
        "introduction": "Imaging Atmospheric Cherenkov Telescopes (IACTs) rely heavily on the atmosphere as their detecting medium. Although the atmosphere gives the telescope systems huge effective areas, daily variations in atmospheric quality can affect the system performance and lead, in the worst cases, to systematic bias in the estimated energy of a given event. Significant effort has been made in the past to take account of this problem by using the cosmic-ray background seen by the telescope on a given night to normalise the data \\cite{bib:bohec}. However, given a better understanding of the location of atmospheric aerosol populations from LIDAR measurements and via modelling of these populations, it is possible to determine an active atmospheric correction to the data. Herein, recent work on such a technique is discussed as applied to observations with the H.E.S.S. telescope array of the active galactic nucleus (AGN) PKS 2155-304. ",
        "conclusions": "A  method for correcting for changes in low-level atmospheric quality is applied to the  variable source PKS 2155-304. The method, based on cosmic-ray trigger-rate, and Ceilometer input, has allowed a corrected set of fluxes for PKS 2155-304 to be produced from data that would otherwise be unusable. This is particularly important as this data set forms part of a large multi-wavelength campaign so removing atmospheric biases is vital. To the lowest order, the effect on integral gamma-ray flux is seen to be proportional to the zenith- and time-corrected cosmic-ray trigger-rate. \\newline A  single scattering LIDAR installed in 2007 at the H.E.S.S. site operate at wavelengths closer to Cherenkov light and has a greater range, and will hopefully allow more straight forward correction. As has been shown, though, the comparison of real and simulated reconstructed shower depth under the application of different atmospheric models allows a coarse appreciation of atmospheric conditions, which is a useful check for the more accurate LIDAR dataset expected to be obtained soon.  A more detailed examination of this technique is presented in \\cite{bib:nolan}.  \\par{\\noindent {\\it Acknowledgements}  The authors would like to acknowledge the support of their host institutions, and in addition we are grateful to Werner Hofmann, Michael Punch, Paula Chadwick \\& Michael Daniel for their helpful comments. S.N. and C.R. acknowledge support from the Science and Technology Facilities Council of the UK. G.P. acknowledges support by the German Ministry for Education and Research (BMBF) through DESY grant 05CH5VH1/0 and DLR grant 50OR0502.}"
    },
    "1402/1402.4765_arXiv.txt": {
        "abstract": "Ultraviolet spectra of stars recorded by orbiting observatories since the 1970\u2019s have revealed absorption features produced by atoms in their favored ionization stages in the neutral ISM of our Galaxy.   Most elements show abundances relative to hydrogen that are below their values in stars, indicating their removal by condensation into solid form.  The relative amounts of these depletions vary from one location to the next, and different elements show varying degrees of depletion.  In a study of abundances along 243 different sight lines reported in more than 100 papers, Jenkins (2009) characterized the systematic patterns for the depletions of 17 different elements, and these results in turn were used to help us understand the compositions of dust grains.  Since the conclusions are based on differential depletions along different sightlines, they are insensitive to errors in the adopted values for the total element abundances.  Some of the more remarkable conclusions to emerge from this study are that (1) oxygen depletions in the denser gas regions (but not as dense as the interiors of molecular clouds) are stronger than what we can expect from just the formation of silicates and metallic oxides, and (2) the chemically inert element krypton shows some evidence for weak depletion, perhaps as a result of trapping within water clathrates or binding with ${\\rm H}_3^+$. ",
        "introduction": "\\label{background}  When various elements condense into solid form in the interstellar medium, there are profound reductions in the relative abundances of many types of heavy atoms and ions in the gas phase. We can obtain insights on the compositions of the dust grains by studying the strengths of depletions of different gas constituents below what we believe to be their intrinsic abundances.  Virtually all of the free atoms in the neutral interstellar medium are in their ground electronic states.  For nearly all of the important atomic species, the energy separations of excited levels above the ground state are larger than the energies of transitions in the visible part of the spectrum.  Hence, as shown in Fig.~\\ref{line_density}, for absorption spectroscopy a vast majority of the useful features are situated in the UV part of the spectrum; this fact was recognized long ago by Spitzer and Zabriskie (1959).  For studying gas in our Galaxy or nearby systems, these transitions can only be accessed by observatories in space.  Except for the elements Be (Spitzer 1949 ; Herbig 1968 ; Boesgaard 1985) and Ti (Wallerstein \\& Goldsmith 1974 ; Stokes \\& Hobbs 1976 ; Stokes 1978 ; Wallerstein \\& Gilroy 1992 ; Welsh et al. 1997 ; Prochaska et al. 2005 ; Ellison et al. 2007 ; Welty \\& Crowther 2010) and some comparisons of the abundances of Ca to Na with velocity (Routly \\& Spitzer 1952 ; Siluk \\& Silk 1974 ; Vallerga et al. 1993 ; Sembach \\& Danks 1994), research on interstellar atomic abundances relied almost entirely on recording stellar spectra over the range extending from the Lyman limit (912$\\,$\\AA) up to about 3000$\\,$\\AA, where the Earth's atmosphere becomes transparent.  \\begin{wrapfigure}{r}{0.5\\textwidth} \\begin{center} \\includegraphics[width=0.48\\textwidth]{line_density002.jpg} \\end{center} \\caption{Line density vs. wavelength\\label{line_density}} \\end{wrapfigure}  The {\\it Copernicus\\/} satellite (Rogerson et al. 1973), launched in 1972, was the first major observatory that had the resolution and sensitivity for productive research on atomic abundances in the interstellar medium (Spitzer \\& Jenkins 1975 ; Cowie \\& Songaila 1986).  The earliest results from this facility (Morton et al. 1973) exhibited a pattern of element depletions that led Field (1974) to conclude that the severities of depletions are strongly linked to the temperatures at which the respective elements condense into solid compounds in matter being expelled from the outer portions of stellar atmospheres.  This formation mechanism, along with the creation of dust in supernova ejecta and the further growth of grains in dense gas complexes in space, still remains a central theme in theories of dust production.  A decade of observations with {\\it Copernicus\\/} reached a climax with the study of element abundances across a broad range of sight lines toward bright stars in our Galaxy (Jenkins et al. 1986), revealing that the strengths of depletions were strongly linked to the average gas densities, indicating that the growth and destruction of grains depend on local conditions.  Following the era of {\\it Copernicus\\/}, major advances in achieving higher accuracy and viewing fainter stars arose from observations by the {\\it Goddard High Resolution Spectrograph\\/} (Brandt et al. 1994) aboard the {\\it Hubble Space Telescope} (HST) soon after its launch in 1990 (Cardelli et al. 1991a ; Cardelli et al. 1991b ; Savage et al. 1991 ; Savage et al. 1992 ; Cardelli et al. 1993 ; Sofia et al. 1993 ; Cardelli 1994).  Savage and Sembach (1996) provided an excellent review of these early accomplishments with HST.  A major upgrade in the ability to record a broad wavelength coverage at high resolution was provided when the {\\it Space Telescope Imaging Spectrograph\\/} was installed on the second servicing mission of HST in 1997 (Woodgate \\& et al. 1998).  The {\\it Far Ultraviolet Spectroscopic Explorer\\/} (FUSE), launched in 1999, extended our coverage to a wavelength band between the Lyman limit and the short- wavelength cutoff of HST.  This facility allowed us to determine the amount of H$_2$ that accompanied the H~I gas and also gave us an access to some important atomic transitions that were not available to HST (Moos et al. 2000). ",
        "conclusions": "A study of the depletions of different elements from the gas phase of the interstellar medium sets some important constraints on the possible compositions of dust grains in different stages of development or levels of concentration. However, we do not obtain explicit information on what compounds are formed. One can, however, make some illustrative guesses on the possible mixes of minerals that satisfy the constraints, as had been done by Draine (2004), who examined the results from a contemporaneous study of depletions by Jenkins (2004).  Some words of caution are in order.  There may be some dangers in over-generalizing the results presented in \\S\\S \\ref{unified}$-$\\ref{buildup} when one considers grain compositions in systems other than our immediate neighborhood in our Galaxy.  One might suppose that in a system where the proportions of different elements are not the same as what we see here, it would be a simple matter to estimate how rapidly the elements disappear into solid form by just integrating the differential depletion relationships, as expressed in Eq.~\\ref{differential_grain_comp}, after starting with some initial gas composition. Complications arise when we recognize that some compounds rely on the presence of others to form initially and remain stable for long periods of time. For instance, Lodders (2003) has pointed out that the incorporation of the elements Ni and Ge in solids depends on the presence of a host element Fe to create an alloy.  Likewise, the formation of refractory compounds that contain Zn and Mn, such as ${\\rm Zn_2SiO_4}$, ${\\rm ZnSiO_3}$, or ${\\rm Mn_2SiO_4}$, are aided by pre-existing host minerals such as forsterite and enstatite.  In some galaxy environment where the $\\alpha/{\\rm Fe}$ differs from that of our own, some elements may have more or less than their respective host compounds that contain other key elements, thus altering the depletion rates in a way that is difficult to predict.  We need also to recognize that any changes in the mix of sources that create the initial grains, such as different kinds of stars and supernovae, can have a profound influence on the outcome.  These initial grains are important constituents by themselves, and moreover they eventually form the cores of grains that undergo further development in the interstellar medium."
    },
    "1402/1402.1681_arXiv.txt": {
        "abstract": "\\textit{It is shown that two $PeV$ events observed by the IceCube collaboration can be interpreted as resonance production of color octet neutrinos.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3151_arXiv.txt": {
        "abstract": "{ HII galaxies are clumpy and their gas kinematics can be mapped to show the global turbulent motions and the effect of massive star evolution. The distribution of their physical conditions is homogenous and oxygen abundance is uniform. The presence of nebular HeII 4868 line seems to be higher in a low abundance galaxy, implying a harder ionization power probably due to stars in low metallicity. Innovative methods of data cube analysis, namely PCA tomography, seem promising in revealing additional information not detected with the standard methods. I review some of our own recent work on the 3D spectroscopy of HII galaxies. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1538_arXiv.txt": {
        "abstract": "To reveal the structures of a transition disk around a young stellar object in Lupus, \\object{Sz\\ 91}, we have performed aperture synthesis 345 GHz continuum and CO(3--2) observations with the Submillimeter Array ($\\sim1\\arcsec$--3$\\arcsec$ resolution), and high-resolution imaging of polarized intensity at the $K_s$-band by using the HiCIAO instrument on the Subaru Telescope ($0\\farcs25$ resolution). Our observations successfully resolved the inner and outer radii of the dust disk to be 65 and 170 AU, respectively, which indicates that \\object{Sz\\ 91} is a transition disk source with one of the largest known inner holes. The model fitting analysis of the spectral energy distribution reveals an H$_2$ mass of $2.4\\times10^{-3}$ $M_\\sun$ in the cold ($T<$30 K) outer part at $65<r<170$ AU by assuming a canonical gas-to-dust mass ratio of 100, although a small amount ($>3\\times10^{-9}$ $M_\\sun$) of hot ($T\\sim$180 K) dust possibly remains inside the inner hole of the disk. The structure of the hot component could be interpreted as either an unresolved self-luminous companion body (not directly detected in our observations) or a narrow ring inside the inner hole. Significant CO(3--2) emission with a velocity gradient along the major axis of the dust disk is concentrated on the \\object{Sz\\ 91} position, suggesting a rotating gas disk with a radius of 420 AU. The \\object{Sz 91} disk is possibly a rare disk in an evolutionary stage immediately after the formation of protoplanets because of the large inner hole and the lower disk mass than other transition disks studied thus far. ",
        "introduction": "Nearly all newly-formed stars are surrounded by disks of gas and dust, which provide the building blocks of planets \\citep{bib:evans2009}. Thus, studying the structure and evolution of these protoplanetary disks provides information about how and when planets form. Initially, the disks are optically-thick, producing broadband infrared (IR) emission well in excess of the stellar photosphere. The star+disk system has a rather flat IR spectral energy distribution (SED) slope \\citep[e.g.,][]{bib:lada1984,bib:lada1987}. By $\\sim5$ Myr, nearly all stars lack evidence of warm circumstellar dust and accretion onto the star \\citep[e.g.,][]{bib:fedele2010,bib:currie2011}. This implies that most protoplanetary disks around solar/subsolar-mass stars have disappeared by this time, and the disk material has been accreted/dispersed because of processes such as photoevaporation or the disk material has been incorporated into planets.\\par  Transition disks bridge the gap between these endpoints of disk evolution because their excess emission at some IR wavelengths is intermediate between those of an optically thick disk and a star lacking a disk \\citep[e.g.,][]{bib:strom1989}. Although the term ``transition disk'' includes a diverse set of morphologies, a major subset of these objects include those with near/mid IR deficits (relative to an optically-thick disk) but optically-thick emission at other, typically longer wavelengths, which is indicative of inner holes or gaps \\citep[e.g.,][]{bib:calvet2002}. IR SED modeling implies that the hole or gap sizes for these objects typically range between $\\sim1$ and $\\sim50$ AU \\citep[e.g.,][]{bib:espaillat2010,bib:merin2010} and encloses the planet-forming region in our own solar system. SED modeling has identified numerous other transition disks in nearby star forming regions \\citep[e.g.,][]{bib:cieza2010,bib:merin2010,bib:currie2011,bib:romero2012,bib:cieza2012b}. Because formation of gas giant planets creates holes or gaps in disks, transition disks with SEDs consistent with these features may be excellent laboratories for studying planet formation.  Although SED modeling can provide indirect evidence of an inner hole or gap, high-resolution observations with long baseline (sub-)millimeter interferometers and large aperture IR telescopes provide direct evidence of their existence \\citep[e.g,][]{bib:brown2009,bib:hughes2009,bib:andrews2011,bib:isella2012,bib:cieza2012a,bib:mathews2012,bib:thalmann2010,bib:hashimoto2012,bib:mayama2012}. Since millimeter and micron sized dust grains contribute to the majority of millimeter and NIR emissions, it is important to compare the detailed spatial distributions of these tracers for understanding the physics of the transition disks.\\par  \\object{Sz 91} is an M0.5 young star surrounded by a transition disk located in the Lupus III molecular cloud \\citep[$d=200$ pc;][]{bib:comeron2008}. The stellar position, ($\\alpha_\\mathrm{J2000}$, $\\delta_\\mathrm{J2000}$), is (16$^\\mathrm{h}$ 07$^\\mathrm{m}$ 11$\\fs$6, $-39\\arcdeg$ 03$\\arcmin$ 47$\\farcs$2). The stellar mass and age have been estimated to be 0.49 $M_\\sun$ and 5 Myr, respectively \\citep{bib:hughes1994}. The H$\\alpha$ line width at 10 \\% of the peak has been measured to be 283 km s$^{-1}$, suggesting accretion with $10^{-10}$ $M_\\sun$ yr$^{-1}$ \\citep{bib:romero2012}. \\object{Sz 91} is classified as a wide binary with a separation of 9\\arcsec \\citep{bib:melo2003}, corresponding to 1800 AU. However, such a wide binary system is unlikely to disturb the disk evolution of a host star \\citep{bib:kraus2012,bib:harris2012}. Moreover the difference between their proper motions implies that these sources are not co-moving \\citep{bib:roeser2010}. We therefore treat the \\object{Sz 91} disk as a circumstellar disk around a single star system throughout this paper.\\par  The SED of \\object{Sz 91} is characteristic of the transition disk. The following remarkable features of this source can be identified in the SED: no significant IR excess with a spectral index from $K_\\mathrm{s}$ to 24 $\\micron$ of $-2$ (i.e., class\\ III in the IR categorization), presence of a large dip of approximately 20 $\\micron$, and very steep flux density rising between 24 and 70 $\\micron$ \\citep{bib:evans2009,bib:romero2012}. \\citet{bib:romero2012} categorized \\object{Sz 91} as a giant planet-forming disk on the basis of the following features: very steep increase in flux density at 24 $\\micron$, which indicates a sharp edge at the inner radius of the disk; a clear sign of mass accretion onto the central star; a relatively massive disk ($\\sim5\\times10^{-3}$ $M_\\sun$).\\par  Although most transition disks with inner holes have flux deficits restricted to 1--10 $\\micron$ and optically thick emission at longer wavelengths, \\object{Sz 91} differs from them because its flux deficit extends to significantly longer wavelengths (24 $\\micron$), and it has strong far-IR to millimeter emission. The total flux density of the 870 $\\micron$ (345 GHz) continuum emission of \\object{Sz 91} has been estimated to be 34.5$\\pm$2.9 mJy \\citep{bib:romero2012}. Although the 1.3 mm continuum emission had not been detected previously \\citep{bib:nurnberger1997}, the recent wide field imaging survey with the AzTEC receiver on the Atacama Submillimeter Telescope Experiment (ASTE) has clearly detected a 1.1 mm flux density of 27.2$\\pm$6.0 mJy \\citep{bib:kawabe2014}. These submillimeter flux densities imply that the \\object{Sz 91}'s disk is substantially massive with respect to those of other class\\ III sources in nearby star forming regions \\citep{bib:andrews2005,bib:andrews2007}; the flux densities roughly correspond to 2--5$\\times$10$^{-3}$ $M_\\sun$ if the canonical gas-to-dust mass ratio of 100 is assumed. Such an object having significant submillimeter flux density, but no NIR excess appears to be at the transition phase from class\\ II to III and is quite rare \\citep[$\\sim$5 \\%][]{bib:andrews2005,bib:andrews2007}. Although a recent high-resolution imaging survey for transition disk objects \\citep{bib:andrews2011} has successfully revealed a large inner hole in the disk up to a radius of $\\sim70$ AU, there is selection bias for the most massive transition disks. Because the submillimeter flux density of \\object{Sz 91} is the lowest among the transition disk objects studied thus far, considering the distance to the source \\citep[Fig.10 in][]{bib:andrews2011}, \\object{Sz 91} is a crucial target for investigating the disk evolution from the perspective of a variety of submillimeter flux densities, i.e., the disk mass. However, the disk structure has not been resolved thus far, and high-resolution imaging observations are urgently required.\\par  In this paper, we present the first sub-arcsecond resolution images toward \\object{Sz 91} at submillimeter and near infrared wavelengths. In section \\ref{sec:observation}, observational parameters and calibration for the obtained data are described. In section \\ref{sec:result}, high-resolution images are shown. In section \\ref{sec:discussion}, the detailed disk structure is discussed on the basis of model fitting to the SED and the CO(3--2) line profile. The unique nature of \\object{Sz 91} is determined from a comparison with other transition disks. Finally in section \\ref{sec:summary}, the key results of this study are summarized.\\par ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Spatially Resolved Disk Structure}\\label{sec:disk_structure} Both the high-resolution submm and NIR images clearly resolved the disk structure around \\object{Sz 91}. In this section, we estimate the inner radius, outer radius, and inclination angle of the disk directly from our images. These parameters are adopted in the model calculations in sections \\ref{sec:sedcalc} and \\ref{sec:linecalc}.\\par  The inner radius of the dust disk can be estimated from the large inner hole revealed in the $K_s$-band PI image. Assuming that the $K_s$-band emission is the scattered light at the inner edge of the disk, the best-fit ellipse in the $K_s$-band image shows the inner radius to be 65$\\pm$4 AU. There are two possibilities to explain the scattered light. One is forward scattering at the surface of the inner part of the disk, and the other is scattering at the inner wall of the disk such as that in the case of \\mbox{LkCa\\ 15} \\citep{bib:thalmann2010}. We can not judge which case is more likely to explain the NIR emission from the Sz 91 disk. However, in either case, it probably originates from the innermost part of the disk, and our estimation of the inner radius seems to be reasonable. The high-resolution 345 GHz image also provides a rough estimate of the disk inner radius. Although the 345 GHz image in Figure \\ref{fig:sma_cont}(b) does not clearly show the presence of the inner hole, it is clear that the 345 GHz flux density becomes weak toward the central star, which is suggestive of dust depletion near the star. The angular distance between the 345 GHz emission peak and the central star can be interpreted as the disk inner radius of $0\\farcs43$ ($=86$ AU), which is roughly consistent with the above estimate of 65 AU. In the following discussion, we adopt the disk radius $R_\\mathrm{in}$ of 65 AU.\\par  The outer radius of the dust disk can be estimated from the beam-deconvolved size of the full-UV 345 GHz map shown in Figure\\ \\ref{fig:sma_cont}(a). Assuming a geometrically thin disk, the beam-deconvolved size of (1\\farcs7$\\pm$0\\farcs1)$\\times$(0\\farcs7$\\pm$0\\farcs2) corresponds to the outer radius $R_\\mathrm{out}$ of 170$\\pm$20 AU at the distance of 200 pc. The gas disk is also resolved in CO(3--2) (Figure\\ \\ref{fig:sma_co_int}). The beam-deconvolved size is measured to be (3\\farcs1$\\pm$0\\farcs3)$\\times$(1\\farcs6$\\pm$0\\farcs7) at a PA of 15$\\pm$11\\degr from a 2-D Gaussian fitting, corresponding to $R_\\mathrm{out}=$310$\\pm$60 AU. However, the red-shifted gas is probably contaminated by the ambient cloud gas, and the center position of the beam-deconvolved disk is shifted by $0\\farcs6$ from the star to the north. Thus we consider the largest distance of $2\\farcs1$ from the stellar position as the outer radius of the gas disk: $R_\\mathrm{out}=$420 AU. Because the actual intensity distribution of each emission is expected to have a power-law form, these outer radii derived from the Gaussian fitting should be regarded as lower limits.\\par   The outer radius of the gas disk is much larger than that of the dust disk. The discrepancy in radial extension between the dust and gas emissions is frequently observed for other protoplanetary disks, which can be explained by an exponential decrement in the surface density \\citep[e.g.,][]{bib:hughes2008,bib:panic2009,bib:isella2010} or by the radial drift of large grains \\citep{bib:andrews2012}. In addition, the possible north-south asymmetry in the dust continuum emission, suggested in Figure \\ref{fig:sma_cont}(b), could have resulted from the azimuthal drift of large grains \\citep{bib:birnstiel2013,bib:vandermarel2013} associated perhaps with a perturbing body orbiting within the central hole of this disk.\\par  The disk inclination angle can be derived from the ratio of the major to minor axis lengths of the best-fit ellipse in Figure \\ref{fig:kband}(a): 40$\\pm$15$\\degr$ from face-on. Moreover, the beam-deconvolved size of the 345 GHz continuum emission in Figure \\ref{fig:sma_cont}(a) provides an inclination angle of 66$\\pm$5$\\degr$. The latter value of 66$\\degr$, however, is possibly influenced by the disk flaring because the distribution of the 345 GHz emission is biased to the outer part of the disk. Thus, we adopt the $K_s$-band value of 40$\\degr$ as the disk inclination angle in this study. In fact, for the CO(3--2) spectrum, the model calculation in section 4 prefers the smaller value of 40$\\degr$.\\par   \\subsection{Disk Parameters Derived from SED fitting}\\label{sec:sedcalc} To deduce other parameters of the dust disk, we performed least-square fitting to the SED on the basis of the power-law disk model. Figure \\ref{fig:sed} shows the SED of \\object{Sz 91} that includes our 345 GHz flux density and a recent data set by {\\it WISE}. As \\citet{bib:romero2012} reported, the SED is characterized by large dip around 20 $\\micron$, sharp rise from 20 to 70 $\\micron$, and significant (sub-)millimeter emission.\\par  We introduced the two components of a cold disk and a small amount of hot dust inside the disk as follows. First, we applied the model for the cold disk, which comprised the usual power-law disk and blackbody stellar emissions, to all of the SED data except that at 10--30 $\\micron$, because the model could not reproduce all the data points including the three points at 10--30 $\\micron$. The power-law disk model has surface density and temperature radial distributions of $\\Sigma(r)$ and $T(r)$, respectively, in a power-law form \\citep{bib:kitamura2002,bib:tsukagoshi2011}: \\begin{equation} \\Sigma(r)=\\Sigma_\\mathrm{in} \\biggl( \\frac{r}{R_\\mathrm{in}} \\biggr)^{-p} \\end{equation} and \\begin{equation} T(r)=T_\\mathrm{in} \\biggl( \\frac{r}{R_\\mathrm{in}} \\biggr)^{-q} \\mathrm{,} \\end{equation} where $r$ is the radial distance, $\\Sigma_\\mathrm{in}$ and $T_\\mathrm{in}$ are the surface density and temperature at the inner radius, $R_\\mathrm{in}$, and $p$ and $q$ are power-law indexes. The gas-to-dust ratio was assumed to be 100 and the extinction of the stellar light because of the interstellar dust was corrected by the $A_\\mathrm{V}$ value toward the star and the dust mass absorption coefficient \\citep[Figure 1 of][]{bib:adams1988}. Table \\ref{tab:sedfit1} lists the fixed parameters in the SED fitting. The stellar mass, the effective temperature of the star, and the visual extinction toward the star are from \\citet{bib:hughes1994}. The inner and outer radii and inclination angle of the disk were determined in section \\ref{sec:disk_structure}. We here adopted the values of 1.5 and 0.5 as $p$ and $q$, respectively, which are the same as those of the minimum mass solar nebular \\citep{bib:hayashi1981}. These values were selected because the SED fitting is known to be insensitive to the power-law index of $p$ and it is difficult to resolve the parameters for the temperature profile ($T$ and $q$) from only the longer wavelength data ($>30$ \\micron). Notably, the lower limit of the disk temperature was set to be 10 K, which is the typical temperature of the Lupus\\ III cloud \\citep{bib:vilas-boas2000}. The stellar radius, $R_\\ast$, $T_\\mathrm{in}$, $\\Sigma_\\mathrm{in}$, and the power-law index of the dust mass opacity coefficient $\\beta$ \\citep[$\\kappa_\\nu=0.1\\times(\\nu/10^{12}\\ \\mathrm{Hz})^\\beta$ cm$^2$ g$^{-1}$:][]{bib:beckwith1990} were treated as the free parameters in the SED fitting.\\par  After the power-law disk model fitting, we introduced the additional {\\it hot component} inside the disk to reproduce the observed flux densities at $\\sim20$ $\\micron$. The presence of the inner hot component is supported by the sign of mass accretion derived from the H$\\alpha$ emission line \\citep{bib:romero2012}. Because no information was available on the structure of this component, we simply assumed a gray body with a temperature of $T_\\mathrm{c}$, a column density of $\\Sigma_\\mathrm{c}$, and a solid angle of $\\Omega_\\mathrm{c}$. Its flux density $S_\\lambda$ at wavelength $\\lambda$ is written by \\begin{equation} S_\\lambda = \\frac{2 h c^2}{\\lambda^5}  \\frac{1}{\\exp(\\frac{hc}{\\lambda k_\\mathrm{B} T_\\mathrm{c}})-1} (1-e^{-\\kappa_\\nu \\Sigma_\\mathrm{c}}) \\times \\Omega_\\mathrm{c}\\ \\mathrm{,} \\end{equation} where $h$ is the Plank constant, $c$ is the speed of light, and $k_\\mathrm{B}$ is the Boltzmann constant. We applied the same dust mass opacity coefficient, $\\kappa_\\nu$, as that in the cold disk. The parameters of $T_\\mathrm{c}$, $\\Sigma_\\mathrm{c}$, and $\\Omega_\\mathrm{c}$ were treated as free. In the fitting, the upper limit of $\\Omega_\\mathrm{c}$ was set to be $\\sim6.0\\times10^{-12}$ str, corresponding to the solid angle of the inner hole seen in the NIR image.\\par  The best-fit model SED of the cold disk reproduced the longer wavelength data ($>30$ \\micron) effectively, as shown in Figure \\ref{fig:sed}. The best-fit parameters of the cold disk are summarized in Table \\ref{tab:sedfit2}. The $\\beta$ value of 0.5$\\pm$0.1 is significantly smaller than that in the diffuse interstellar medium of $\\sim2$ \\citep{bib:draine1984}, suggesting that the dust growth occurs in the disk \\citep{bib:miyake1993}. By adopting this $\\beta$ and $\\kappa_\\nu$ introduced by \\citet{bib:beckwith1990}, who assumed a 100:1 mass ratio between gas and dust, the disk mass is derived to be (2.4$\\pm$0.8)$\\times10^{-3}$ $M_\\sun$. This value is significantly higher than those of other class III sources in nearby star forming regions \\citep{bib:andrews2005,bib:andrews2007}; however, the value is lower than those of most of the transition disks studied thus far \\citep{bib:andrews2011}, even if we consider the difference in $\\kappa_\\nu$ by a factor of $\\sim$1.7 at $\\nu=340$ GHz between this study and the previous studies in which $\\beta$ is set to be 1.\\par  To verify the fitting result, we also created a continuum image for the best-fit model and compared it with the observations, as shown in Figure \\ref{fig:modelim}. The model image essentially agrees with the observations, but there remains a difference of at most $4\\sigma$; the negative residual extends in the east-west direction, which may be due to the asymmetry of the disk.\\par  We next attempted to reproduce the observed flux densities at 10--30 $\\micron$ by adding the contribution from the hot component. The best-fit parameters of $T_\\mathrm{c}$, $\\Sigma_\\mathrm{c}$, and $\\Omega_\\mathrm{c}$ to reproduce all of the SED data were searched in the reduced $\\chi^2$ maps by manually changing the initial values of $T_\\mathrm{c}$, $\\Sigma_\\mathrm{c}$, and $\\Omega_\\mathrm{c}$ with 6000 runs.\\par  From the calculations, we determined that there are two distinct regions in the $\\Sigma_\\mathrm{c}$-$\\Omega_\\mathrm{c}$ plane where the reduced $\\chi^2$ takes its local minimum values, which indicates a strong coupling between $\\Sigma_\\mathrm{c}$ and $\\Omega_\\mathrm{c}$. One is the region in which $\\Sigma_\\mathrm{c}<5\\times10^{-5}$ g cm$^{-2}$ and $\\Omega_\\mathrm{c}>4\\times10^{-13}$ str, and the emission of the hot component is substantially optically thin. In this region, the best-fit $T_\\mathrm{c}$ typically converges to be 186 K and the best-fit $\\Sigma_\\mathrm{c}$ is inversely proportional to $\\Omega_\\mathrm{c}$. The product of $\\Sigma_\\mathrm{c}\\times\\Omega_\\mathrm{c}$ is a constant value, which provides the best-fit mass of $3\\times10^{-9}$ $M_\\sun$ for the hot component. The other is the region in which $\\Sigma_\\mathrm{c}>30$ g cm$^{-2}$ and the emission is substantially opticall thick. In this region, the best-fit $T_\\mathrm{c}$ is typically 172 K and $\\Sigma_\\mathrm{c}$ is independent of $\\Omega_\\mathrm{c}$. For the hot component, we can obtain the best-fit $\\Omega_\\mathrm{c}=1\\times10^{-16}$ str and the lower limit mass of $6\\times10^{-7}$ $M_\\sun$. Although the reduced $\\chi^2$ values are slightly lower in the former case ($\\Delta\\chi^2=0.3$), the difference between the two regions is not significant. We therefore conclude that there are two possible origins of the hot component: the optically thin gray body emission and the optically thick black-body emission. However, the optically thin condition is the unlikely case because the mass of the hot component is too small given the mass accretion rate of $10^{-10}$ $M_\\sun$ yr$^{-1}$; all the hot component would disappear in only 30 yrs.\\par  Notably, the contribution of the hot component is mainly restricted by only three data points at MIR, from 12 to 25 $\\micron$; the upper limit of the temperature is determined by the slight excess emission at 12 $\\micron$ and the lower limit is limited by the decrement between 22 and 24 $\\micron$. A refined model will be required when more data at MIR and FIR are obtained by further observations.\\par    \\begin{figure} \\epsscale{1.0} \\plotone{figure5.eps} \\caption{Spectral energy distribution of \\mbox{Sz 91}. The filled circles indicate observed flux densities. All data except the 345 GHz flux density were compiled from previous studies, including the NOMAD catalog \\citep{bib:zacharias2005}, the {\\it 2MASS} point source catalog \\citep{bib:cutri2003}, the {\\it Spitzer} IRAC and MIPS photometry \\citep{bib:evans2009}, the {\\it WISE} all-sky data release \\citep{bib:cutri2012}, the {\\it AKARI} FIS all-sky survey point source catalog \\citep{bib:yamamura2010}, the {\\it IRAS} point source reject catalog \\citep{bib:iras2007}, the AzTEC/ASTE 1100 $\\micron$ photometry \\citep{bib:kawabe2014}, and the {\\it Herschel} PACS/SPIRE FIR flux densities which we derived from the HSA Science archive data. The best-fit SED and the contribution of the hot component in the optically thin case are shown by the black and red lines, respectively, and by the gray and orange lines in the optically thick case. The blue dashed line indicates the stellar contribution.}\\label{fig:sed} \\end{figure}  \\begin{figure} \\epsscale{1.0} \\plotone{figure6.eps} \\caption{Comparison of the observed data with the best-fit model. The 345 GHz continuum image, model image, and residual image (data$-$model) are shown from left to right. The contour starts at $\\pm2\\sigma$, where 1$\\sigma=$2.1 mJy beam$^{-1}$. The solid and dashed lines indicate positive and negative, respectively.}\\label{fig:modelim} \\end{figure}  \\begin{deluxetable}{cc} \\tabletypesize{\\small} \\tablecaption{Fixed parameters in the SED fitting} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter} & \\colhead{Fixed Value} } \\startdata $M_\\ast$ ($M_\\sun$) & 0.49\\tablenotemark{1} \\\\ $A_\\mathrm{V}$ (mag) & 2.0\\tablenotemark{1} \\\\ $R_\\mathrm{in}$ (AU) & 65 \\\\ $R_\\mathrm{out}$ (AU) & 170 \\\\ $i$ (degree) & 40 \\\\ $p$ & 1.5 \\\\ $q$ & 0.5 \\\\ \\enddata \\label{tab:sedfit1} \\tablenotetext{1}{\\citet{bib:hughes1994}} \\end{deluxetable}  \\begin{deluxetable}{cc} \\tabletypesize{\\small} \\tablecaption{Best-fit parameters in the SED fitting} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter} & \\colhead{Value} } \\startdata \\multicolumn{2}{c}{Stellar parameters}\\\\ \\hline $R_\\ast$ ($R_\\sun$) & 1.29$\\pm$0.03 \\\\ $T_\\ast$ (K) & 4148$\\pm$55 \\\\ $L_\\ast$ ($L_\\sun$) & 0.49$\\pm$0.02 \\\\ \\hline \\multicolumn{2}{c}{Cold outer disk} \\\\ \\hline $T_\\mathrm{in}$ (K) & 32.5$\\pm$3.9 \\\\ $\\Sigma_\\mathrm{in}$\\tablenotemark{1} (g cm$^{-2}$) & 0.67$\\pm$0.03 \\\\ $\\beta$ & 0.5$\\pm$0.1 \\\\ $M_\\mathrm{disk}$\\tablenotemark{1} (10$^{-3}$ $M_\\sun$) & 2.4$\\pm$0.8 \\\\ \\enddata \\label{tab:sedfit2} \\tablenotetext{1}{The surface density and the disk mass are shown in the gas+dust density and mass by assuming the gas-to-dust mass ratio of 100.} \\end{deluxetable}   \\subsection{Origin of the Hot Component in the Inner Hole of the Disk}\\label{sec:sedimp} The SED model analysis indicates that the presence of a hot component inside the dust disk with a temperature of $\\sim$180 K, which is responsible for the SED peak at $\\sim$20 $\\micron$. There are two possible origins for such a hot component: a localized self-luminous emitting body (i.e., circumplanetary disk) or an inner warm structure of the disk \\citep{bib:wolf2005}. In this section, we estimate the size and mass of the hot component for both cases.\\par  In the case of the circumplanetary disk, we can estimate the radius from $\\Omega_\\mathrm{c}$ because it is usually optically thick \\citep{bib:dangelo2003}. Assuming the circumplanetary disk is parallel to the parent circumstellar disk, the solid angle of $1\\times10^{-16}$ str corresponds to the radius of 0.3 AU, or 64 $R_\\sun$. On the other hand, the Hill radius of a putative planet around \\object{Sz 91} is expressed by \\begin{equation} R_\\mathrm{Hill} = 1.1 \\times \\frac{a}{10\\ \\mathrm{AU}} \\biggl( \\frac{m}{1\\ M_\\mathrm{J}}\\biggr)^{1/3} \\ \\mathrm{[AU],} \\end{equation} where $m$ and $a$ are the mass and orbital radius of the planet, respectively. The radius of the circumplanetary disk inferred from the SED fitting is significantly smaller than $R_\\mathrm{Hill}$ at 3--65 AU from the star, which is consistent with the theoretical expectation for a circumplanetary disk \\citep[e.g.,][]{bib:tanigawa2012}. However, our ADI observation with the Subaru telescope could not examine the presence of such a companion planet due to a low rotation angle of $\\sim16\\degr$.\\par  For the second possibility, because the hot component could be fitted with a single temperature, it must be confined to a narrow width in radius. If we extrapolate the temperature distribution of the outer disk determined in the SED fitting, the temperature of the hot component in the optically thick case (172 K) corresponds to a radius of 2.3 AU. The best-fit $\\Omega_\\mathrm{c}=1\\times10^{-16}$ str corresponds to a ring width of 0.01 AU, which is significantly narrow with respect to the ring radius. In contrast, the optically thin condition is unlikely because the solid angle is $>4\\times10^{-13}$ str which corresponds to a ring width of $>$53 AU. Therefore we conclude that the other implication of the hot component is the optically thick ring at 2.3 AU, whose total mass is at least $6\\times10^{-7}$ $M_\\sun$. Such an example of the optically thick ring around the transition disk has also been reported in \\mbox{RX J1633.9-2442} \\citep{bib:cieza2012a}. The inner structure of \\object{Sz 91} may be similar to that of \\mbox{RX J1633.9-2442}, whereas \\object{Sz 91} exhibits a larger inner hole and a lower disk mass.\\par   \\subsection{Velocity Structure of the Gas Disk: Model Calculation and Comparison with CO(3--2) Profile}\\label{sec:linecalc} The CO(3--2) image shown in Figure \\ref{fig:sma_co_int} suggests the presence of a rotating gas disk around the star. To reveal the disk rotation in detail, we calculated model spectra of CO(3--2) with a simple power-law disk model according to that reported by \\citet{bib:kitamura1993}, and we compared the results with the observed CO(3--2) profile.  The following parameters of the model disk were estimated from the observed images and the SED fitting: the inner and outer disk radii of 65 and 420 AU, respectively; the temperature distribution of $T(r)=32.5\\times(r/R_\\mathrm{in})^{-0.5}$ K, and the surface density distribution of $\\Sigma(r)=0.67\\times(r/R_\\mathrm{in})^{-1.5}$ g cm$^{-2}$. The hydrostatic equilibrium is assumed along the vertical direction and the density distribution, $\\rho(r,z)$, is therefore expressed by \\begin{equation} \\rho(r,z) = \\rho(r,0) \\exp \\biggl[ - \\frac{1}{2} \\biggl( \\frac{z}{H(r)} \\biggr)^2 \\biggr] \\mathrm{,} \\end{equation} where $H(r)$ is the scale height given by \\begin{equation} H(r) = \\sqrt{\\frac{r^3 k_\\mathrm{B}T(r)}{GM_\\ast \\mu m_\\mathrm{H}}} \\mathrm{.} \\end{equation} Here, $G$ is the gravitational constant, $M_\\ast$ is the stellar mass, $\\mu$ is mean molecular weight, and $m_\\mathrm{H}$ is the mass of H atom. The density at the midplane is expressed by \\begin{equation} \\rho(r,0)=\\frac{\\Sigma(r)}{\\sqrt{2\\pi}H(r)} \\mathrm{.} \\end{equation} The gas motion is assumed to be Kepler rotation whose velocity field is written by \\begin{equation} V(r) = \\biggl( \\frac{G M_\\ast}{r} \\biggr)^{0.5} \\mathrm{.} \\end{equation} The fractional abundance of CO with respect to H$_2$ is assumed to be $9\\times10^{-5}$, which corresponds to a typical interstellar value \\citep[e.g.,][]{bib:schloerb1984,bib:irvine1985}. Because the density of the disk is at least $10^6$ cm$^{-3}$ at the midplane which is significantly higher than the critical density of CO(3--2) ($\\sim10^4$ cm$^{-3}$), we assume local thermodynamic equilibrium (LTE). Notably, the hot component in the dust hole was not included because its contribution was negligible ($\\lesssim1\\times10^{-4}$ Jy).\\par  Figure \\ref{fig:sma_co_spec} shows the calculated CO spectra superimposed on the observed CO spectrum integrated over a $5\\arcsec\\times5\\arcsec$ box centered at the stellar position. We noted that the peak intensity of the calculated CO spectrum differed from the observed value. However, the discrepancy is not significant because the disk emission around the stellar LSR velocity is probably contaminated by the ambient cloud emission. Therefore, we focused on the emission at the blue-shifted side ($V_\\mathrm{LSR}<2.9$ km s$^{-1}$). The model profile in the $i=40\\degr$ case agrees with the observed data, confirming the validity of the disk parameters derived from the dust disk. Notably, the inclination angle of 66$\\degr$ determined by the 345 GHz image did not fit well the CO line shape, indicating that the inclination of the dust disk estimated from the NIR scattered light is more plausible.  The presence of the gas inside the inner edge of the dust disk could not be confirmed from our data set. Although we calculated the model spectrum for $R_\\mathrm{in}=$0 AU as shown in Figure \\ref{fig:sma_co_spec}, the difference is within the $1\\sigma$ uncertainty. The presence of the gas disk in the inner hole is supported by the fact that \\object{Sz 91} shows the mass accretion onto the star, and thus, higher spatial resolution imaging is required to reveal the inner structure of the gas disk.\\par  \\begin{figure} \\epsscale{1.0} \\plotone{figure7.eps} \\caption{CO(3--2) spectrum integrated over a $5\\arcsec \\times 5\\arcsec$ area centered at the stellar position (black line). The red line shows the calculated profile by a disk model with an inner radius of 65 AU and an inclination angle of 40$\\degr$ from face-on (see text for the detailed description of the model). The blue line indicates the case of the inner radius of 0 AU. The green line indicates the model profile of the disk with a 65 AU inner radius and an inclination of 66$\\degr$. The CO emission in the velocity range from 4 to 7 km s$^{-1}$ in $V_\\mathrm{LSR}$ is probably contaminated by the ambient cloud emission, as shown by the gray line at the top of the panel.}\\label{fig:sma_co_spec} \\end{figure}    We present the results of the aperture synthesis 345 GHz continuum and CO(3--2) line emission observations with SMA toward a transition disk object in Lupus, \\object{Sz 91}. Furthermore, a high resolution image of the polarized intensity at the $K_s$ band obtained with the Subaru Telescope is also presented. The transition disk around \\object{Sz 91} has been directly resolved and imaged in this study. The disk parameters are derived and the structure and evolutional phase of the \\object{Sz 91} transition disk are discussed. The main results of our observations are summarized in the following points:\\par  \\begin{itemize} \\item{Our high-resolution imaging revealed a dust disk around \\object{Sz 91} with inner and outer radii of 65 and 170 AU, respectively, and an inclination angle of 40$\\degr$. Furthermore, the Kepler rotating gas disk with a radius of 420 AU was imaged in the CO(3--2) line.} \\item{Model analysis of the SED of \\object{Sz 91} was performed by using a simple power-law disk model. We determined that the observed SED can be reproduced well by the combination of a cold disk and a hot component in the inner hole of the disk. The total H$_2$ mass of the cold disk is estimated to be $2.4\\times10^{-3}$ $M_\\sun$ if the canonical gas-to-dust mass ratio of 100 is adopted. The disk mass is significantly higher than those of other class III sources in nearby star forming regions; however,the disk mass is one of the lowest masses among the currently known transition disks.} \\item{We determined that the hot component can be expressed by a single temperature gray body of $\\sim180$ K. Although the hot component could not be resolved by our observations, its origin is either a localized self-luminous emitting body (i.e., a Jovian mass protoplanet with a circumplanetary disk) or an optically thick ring in the inner hole of the disk at 2.3 AU.} \\item{Our results confirm the previous results such that the disk structure of \\object{Sz 91} is consistent with that of an ongoing giant planet forming disk. In particular, the relatively large inner hole and lower disk mass indicate that the transition disk of \\object{Sz 91} is probably in a stage of nearly completing planet formation. \\object{Sz 91} will be a crucial target for investigating the evolution of transition disks and the planetary formation process. In the near future, our proposed study with the Atacama Large Millimeter/submillimeter Array (ALMA) will provide a new insight into the planet formation process.} \\end{itemize}"
    },
    "1402/1402.0014_arXiv.txt": {
        "abstract": "\\noindent We use particle-in-cell (PIC) simulations to study the nonlinear evolution of ion velocity space instabilities in an idealized problem in which a background velocity shear continuously amplifies the magnetic field. We simulate the astrophysically relevant regime where the shear timescale is long compared to the ion cyclotron period, and the plasma beta is $\\beta \\sim 1-100$. The background field amplification in our calculation is meant to mimic processes such as turbulent fluctuations or MHD-scale instabilities.  The field amplification continuously drives a pressure anisotropy with $p_\\perp > p_\\parallel$ and the plasma becomes unstable to the mirror and ion cyclotron instabilities.  In all cases, the nonlinear state is dominated by the mirror instability, not the ion cyclotron instability, and the plasma pressure anisotropy saturates near the threshold for the linear mirror instability. The magnetic field fluctuations initially undergo exponential growth but saturate in a secular phase  in which the fluctuations grow on the same timescale as the background magnetic field (with $\\delta B \\sim 0.3 \\langle B \\rangle$ in the secular phase). At early times, the ion magnetic moment is well-conserved but once the fluctuation amplitudes exceed $\\delta B \\sim 0.1 \\, \\langle B \\rangle$, the magnetic moment is no longer conserved but instead changes on a timescale comparable to that of the mean magnetic field. We discuss the implications of our results for low-collisionality astrophysical plasmas, including the near-Earth solar wind and low-luminosity accretion disks around black holes. \\newline ",
        "introduction": "\\label{sec:intro}  \\noindent Ion pressure anisotropies are ubiquitous in heliospheric and astrophysical plasmas. In the absence of Coulomb collisions, the magnetic moment of ions, $\\mu_i$ ($\\equiv v_{\\perp,i}^2/B$, where $v_{\\perp,i}$ is the ion velocity perpendicular to the local magnetic field $\\vec{B}$ and $B=|\\vec{B}|$) is an adiabatic invariant. Thus, if ion collisions are infrequent, the ion velocity distributions parallel and perpendicular to the magnetic field decouple, making $\\Delta p_i \\equiv p_{\\perp,i} - p_{||,i} \\ne 0$ (where $p_{\\perp,i}$ and $p_{||,i}$ are the pressure components of ions perpendicular and parallel to $\\vec{B}$). Examples of systems where ion pressure anisotropies are important are low-luminosity accretion disks around compact objects \\citep{SharmaEtAl06}, the intracluster medium (ICM) \\citep{SchekochihinEtAl05, Lyutikov07}, and the heliosphere \\citep[see, e.g., ][]{hellinger2006, MatteiniEtAl07, MarucaEtAl11}. In these systems the ion pressure anisotropy is expected to play a key role in the large scale dynamics of the plasma.  For instance, \\cite{SharmaEtAl06} pointed out that pressure anisotropies can give rise to an anisotropic stress that can contribute to angular momentum transport and plasma heating in low-luminosity accretion disks. This physics is not included in standard MHD models of accretion disks.\\newline  \\noindent Velocity space instabilities limit the growth of the pressure anisotropy and give rise to small scale fluctuations in the magnetic field. These fluctuations can, in turn, affect the large scale transport properties of the plasma by modifying the mean free path of particles.\\newline  \\noindent Magnetic field amplification in a collisionless system generically drives $p_{\\perp,i} > p_{||,i}$. There are two ion-scale instabilities that can be excited in this regime: the mirror and the ion-cyclotron (IC) instabilities \\citep{Hasegawa69, Gary92,Southwood93}. The mirror instability consists of non-propagating, strongly compressional modes. Their fastest growing wave vectors $\\vec{k}$ are oblique to $\\vec{B}$, with the magnetic variations parallel to $\\vec{B}$ being much larger than the perpendicular fluctuations, $\\delta \\vec{B}_{||} \\gg \\delta \\vec{B}_{\\perp}$. For a bi-Maxwellian distribution of ions, and assuming cold electrons, the threshold condition for mirror instability growth is given by $T_{\\perp,i}/T_{||,i}-1 > 1/\\beta_{\\perp,i}$, where $T_{\\perp,i}$ ($T_{||,i}$) is the ion temperature perpendicular (parallel) to $\\vec{B}$, and $\\beta_{\\perp,i} \\equiv 8\\pi p_{\\perp,i}/B^2$ \\citep{Hasegawa69}. The IC instability, on the other hand, consists of transverse electromagnetic waves, with the fastest growing $\\vec{k}$ being preferentially parallel to $\\vec{B}$ \\cite[e.g.,][]{AndersonEtAl91}. Whether the IC or mirror instability sets in first depends on how fast these instabilities grow for a given set of plasma conditions. % \\cite{GaryEtAl94} showed that for $T_{\\perp,i}/T_{||,i}-1 = 0.35/\\beta_{||,i}^{0.42}$, the growth rate of the IC instability is $\\gamma_{IC} = 10^{-4}\\omega_{c,i}$, where $\\omega_{c,i}$ is the cyclotron frequency of the ions. In addition, for $\\gamma_{IC}/\\omega_{c,i} \\ll 1$, the threshold anisotropy depends very weakly on $\\gamma_{IC}$. \\newline  \\noindent In the linear regime, the dominant instability will be determined by which threshold is reached first as $\\vec{B}$ is amplified. These estimates for the threshold conditions imply that, in the linear regime, the IC instability should dominate for $\\beta_i \\sim 1$, while the mirror instability should dominate for $\\beta_i \\gg 1$.\\newline  \\noindent Although both the mirror and IC instabilities have been extensively studied in the linear regime, a complete theory of their nonlinear evolution and saturation is still lacking. In this paper, we are particularly interested in the question of how the mirror and IC modes behave after the initial phase of exponential growth. Indeed, in most astrophysical scenarios where the magnetic field is amplified, the growth of the field occurs on time scales much longer than the growth time of the relevant kinetic instabilities. Therefore, most of the evolution of the velocity space instabilities happens in a nonlinear regime, where the conditions typically assumed in linear studies, such as a homogeneous plasma or a bi-Maxwellian distribution of particle velocities, are not necessarily satisfied. Moreover, if the field is amplified by order unity or more, a quasi-linear analysis is not applicable. One of the key questions we are interested in addressing in this case is whether the mirror instability saturates via pitch-angle scattering that violates magnetic moment conservation \\citep{SharmaEtAl06} or via a nearly $\\mu_i$-conserving rearrangement of the magnetic field, with fluctuations $|\\delta \\vec{B}| \\sim |\\vec{B}|$ \\citep{SchekochihinEtAl08}.  We shall see that both of these saturation processes can in fact be important.  \\newline  \\noindent The nonlinear regime of the IC and mirror instabilities has been studied both theoretically \\citep[e.g.,][]{SchekochihinEtAl08,HellingerEtAl09} and numerically  \\citep{BaumgartelEtAl03,TravnicekEtAl07,CalifanoEtAl08, GuoEtAl2009}. In most numerical studies, however, the pressure anisotropy is treated as an initial condition, which decreases as the instabilities grow and saturate. Thus, in these approaches the nonlinear saturation cannot be followed for a time much longer than that of the initial exponential growth.  Moreover, many of the previous numerical studies have focused on one-dimensional simulations.  \\newline  \\noindent In this paper we study the long term, nonlinear evolution of the mirror and IC instabilities using two-dimensional particle-in-cell (PIC) simulations in driven systems. In order to self-consistently explore a time much longer than the initial exponential growth, we continuously induce the growth of $\\Delta p_i$ by amplifying the mean field $<\\vec{B}>$ during the simulation. We concentrate on amplification by incompressible plasma motions. Thus,  our simulations impose a shear velocity on the plasma, which sustains the growth of the magnetic field and maintains an overall positive pressure anisotropy. This growth is intended to mimic magnetic field fluctuations in a turbulent plasma and MHD-scale instabilities that amplify the magnetic field, like the magnetorotational instability \\citep[MRI;][]{BalbusEtAl91} or convective instabilities driven by anisotropic thermal conduction in dilute plasmas \\citep{Balbus01,Quataert08}.  We suspect that our results are also relevant to cases in which perpendicular heating or parallel cooling drives a plasma mirror and/or IC unstable, as in, e.g., the solar wind.  \\cite{HellingerEtAl08} conducted an analogous study to ours using hybrid simulations. In their study, expansion of the simulation box decreased the mean magnetic field, so that on average $p_{\\perp,i} < p_{||,i}$. This in turn led to the growth and nonlinear saturation of the firehose instability.  Also, \\cite{TravnicekEtAl07} conducted a similar hybrid simulation study where the simulation box is expanded along the directions parallel and perpendicular to $\\vec{B}$. With this setup, $p_{\\perp,i}$ becomes larger than $p_{||,i}$ on average, and both the mirror and IC instabilities can grow and reach saturation. Our work is complementary to these hybrid simulation studies and focuses on the process in which the condition $p_{\\perp,i} > p_{||,i}$ is achieved by field amplification due to shearing plasmas. We also consider higher values of $\\beta_i$ $\\approx 1-80$, relevant to accretion disks, the ICM, and the near-Earth solar wind. Our calculations are less directly applicable to heliospheric measurements of mirror modes driven unstable by rapid heating through the Earth's bow shock \\citep[e.g.,][]{SchwartzEtAl96}. This is because we assume that the pressure anisotropy is generated on a timescale long compared to the ion cyclotron timescale, which is probably not true at collisionless shocks.\\newline  \\noindent Our paper is organized as follows. In \\S \\ref{sec:numsetup} we describe the simulation set up, emphasizing the key physical and numerical parameters. \\S \\ref{sec:results} shows our results and \\S \\ref{sec:conclu} presents our conclusions.  Throughout the paper we compare some of our simulation results with linear theory predictions appropriate to our PIC simulations.  These have artificially low ion to electron mass ratios ($m_i/m_e \\simeq 1-10)$.  The linear theory results are based on the linear Vlasov solver developed in \\citet{Verscharen2013}.\\newline  \\noindent During the completion of this work, \\citet{KunzEtAl14} presented calculations of firehose and mirror saturation in shearing plasmas with very similar results to those that we present here.\\newline ",
        "conclusions": "\\label{sec:conclu}  \\noindent We have used PIC simulations to study the nonlinear, long-term evolution of ion velocity space instabilities in collisionless plasmas. We have focused on instabilities driven by pressure anisotropy with $p_{\\perp,i} > p_{\\parallel,i}$:  the ion cyclotron (IC) and mirror instabilities. These instabilities are expected to arise when turbulence and/or MHD instabilities amplify $\\vec{B}$, rendering $p_{\\perp,i} > p_{\\parallel,i}$ due to conservation of the ion magnetic moment $\\mu_i$.  Alternatively, in the solar wind $p_{\\perp,i}/p_{\\parallel,i}$ can increase due to perpendicular heating and/or parallel cooling of the plasma.\\newline  \\noindent In contrast to previous studies, we do not consider the initial value problem of the evolution of a given initial ion pressure anisotropy. Instead, we self-consistently induce the growth of an ion pressure anisotropy by continuously amplifying the mean magnetic field in our computational domain. This setup allows us to study the long-term, saturated state of the plasma, rather than the (much shorter) initial, exponential stage.  The former is of most interest for the heliospheric and astrophysical applications of our work. We have focused on the regime $\\beta_i \\approx 1-80$ and plasma magnetizations $\\omega_{c,i}/s$ in the range $\\omega_{c,i}/s \\sim 100-1000$ (the magnetization is the ratio of the initial ion cyclotron frequency to the background shear frequency).  Most of our simulations used $m_i/m_e=1$, but we showed that our results for the ion scale physics are fairly independent of the mass ratio (see \\S \\ref{sec:compmime}).  So long as $\\omega_{c,i}/s \\gg 1$, our results are also relatively independent of the precise value of the magnetization (see, in particular, Figs. \\ref{fig:beta20B2s} \\& \\ref{fig:mubeta20}). Thus we believe that our basic conclusions about the saturation of the mirror and IC instabilities in continuously driven systems can be applied to physical problems where $\\omega_{c,i}$ is many orders of magnitude larger than $s$.\\newline  \\noindent Our primary results are as follows: \\begin{itemize} \\item For the $\\beta_i \\approx 1-80$ regime we have studied, the turbulent fluctuations in the saturated state are dominated by mirror modes. This is despite the fact that IC and mirror instabilities are both present at early phases (and have similar linear instability thresholds) when $\\beta_i \\approx$ {a few}.  A plausible explanation for this dominance of the mirror instability in the saturated state even at $\\beta_i \\sim$ a few is that the ion velocity distribution departs from a bi-Maxwellian spectrum in the presence of the IC instability.    This can in turn increase the threshold for the IC instability \\citep{Isenberg12,Isenberg13} making it less important than the mirror instability in the nonlinear regime. \\footnote{The mirror instability can dominate the dynamics even more when the anisotropic pressure of the electrons \\citep{RemyaEtAl13} and the full 3D geometry of the problem \\citep{ShojiEtAl09} are taken into consideration.}  We indeed see more significant departures from a bi-Maxwellian in the low $\\beta_i$ regime where the IC instability is present (Figs. \\ref{fig:spectrabeta6} \\& \\ref{fig:spectrabeta20}).   By contrast, the ion distribution function remains relatively bi-Maxwellian in the high $\\beta_i$ regime in which the mirror instability dominates over the IC instability at all times (Fig. \\ref{fig:spectrabeta20}). The dominance of the mirror instability at $\\beta_i \\approx$ a few is consistent with the nonlinear study of \\cite{TravnicekEtAl07}, which makes use of hybrid simulations of an expanding box.  \\item Small-scale fluctuations driven by the mirror instability initially grow exponentially but saturate in a secular phase in which $|\\delta \\vec{B}| \\sim 0.2-0.3 |<\\vec{B}>|$ (Figures \\ref{fig:beta6} \\& \\ref{fig:beta20B2s}).  The fluctuations in the saturated state have $k_{\\perp} \\gg k_{||}$ (see, e.g., Figure \\ref{fig:fieldsbeta6}) consistent with linear theory expectations for mirror modes.  \\item The ion pressure anisotropy $\\Delta p_i/p_{||,i}$ in the saturated state is reasonably well described by the marginal stability condition for the mirror instability (e.g., Figs \\ref{fig:beta6} \\& \\ref{fig:beta20B2s}).   Note, however, that the marginal stability state of the mirror mode in our simulations is not the same as in real systems because of the comparable ion and electron masses and the smaller $\\omega_{c,i}/s$ used in our calculations. This is important to bear in mind when comparing our results to heliospheric measurements.  \\item The total thermal energy of the plasma increases in time at the theoretically predicted rate given the background shear in the plasma and the anisotropic stress associated with the pressure anisotropy (Fig. \\ref{fig:anistress}).  This is consistent with the idea that the ``viscous'' heating rate in a collisionless plasma is set by how velocity space instabilities regulate the pressure anisotropy \\citep{SharmaEtAl06}.   \\item When the growth time of the mirror instability is long compared to the ion cyclotron period -- as is the case in our simulations where the shear slowly amplifies the mean magnetic field -- the ion magnetic moment is initially reasonably well conserved. This is fundamentally because there is no mechanism to violate magnetic moment conservation during the linear phase of mirror growth. So the modes continue to grow to larger and larger amplitudes, given the sustained free energy source created by the background velocity shear amplifying the magnetic field and continuously generating pressure anisotropy. This regime lasts until $\\delta B \\sim 0.1 \\, \\langle B \\rangle$ and $ts \\lesssim 1$, and is consistent with the ``secular regime\"  found by \\citet{KunzEtAl14} where $\\mu$ is nearly conserved.  \\item Saturation of the mirror modes starts to happen once the fluctuations attain $\\delta B \\sim 0.1 \\, \\langle B \\rangle$. In this stage the variation in magnetic field on the scale of the larmor radius is sufficiently large that the ion magnetic moment is no longer conserved. This is consistent with theories of stochastic ion motion in large amplitude turbulent fluctuations \\citep[e.g.,][]{ChenEtAl01, JohnsonEtAl01}. In this phase the ion magnetic moment changes on nearly the same timescale as the mean magnetic field (Fig. \\ref{fig:muderivs}), which is consistent with pitch angle scattering maintaining marginal stability. In this stage the magnetic fluctuations only experience a modest additional growth, reaching a maximum amplitude of $\\delta B \\sim 0.3 \\, \\langle B \\rangle$.  \\item The need for finite amplitude fluctuations with $\\delta B \\sim 0.1 \\, \\langle B \\rangle$ in order to stop magnetic moment conservation implies the existence of a the temporal delay of $\\sim 0.3 s^{-1}$ between the onset of the IC and mirror instabilities and the onset of efficient pitch angle scattering ($s$ is the rate at which the mean magnetic field is growing). \\end{itemize}  \\noindent Our results have a number of consequences for modeling the large scale dynamics of nearly collisionless astrophysical plasmas and for interpreting heliospheric measurements.  We briefly mention a few of these applications but defer a more detailed analysis to future work.\\newline  \\noindent Our results are consistent with a number of in situ heliospheric measurements. For instance, our maximum saturated value of $\\delta B/B \\sim 0.3$ is in reasonable agreement with mirror modes observations in different environments like planetary magnetosheaths \\citep[where large amplitude mirror modes are frequently found. See, e.g.,][]{TsurutaniEtAl82, JoyEtAl06, SoucekEtAl08, VolwekEtAl08, HorburyEtAl09,TsurutaniEtAl11a}, the solar wind \\citep[e.g.,][]{LiuEtAl06, bale2009, RussellEtAl09, Enriquez-riveraEtAl10}, and the heliosheath \\citep[e.g.,][]{TsurutaniEtAl11a,TsurutaniEtAl11b}.\\footnote{We note that our quoted maximum $\\delta B/B\\sim 0.3$ corresponds to $<(\\vec{B}-< \\vec{B}>)^2>^{1/2} / |<\\vec{B}>|$, where $< >$ stands for volume average. This means that our results still allow for {\\it local} magnetic fluctuations of even larger amplitude. It is also important to stress that if a plasma is rapidly driven to have a temperature anisotropy well in excess of the mirror threshold \\citep[e.g., at the Earth's bow shock;][]{TsurutaniEtAl82} then the resulting amplitude of the mirror modes could significantly exceed what we find in our calculations.    In this case, a standard initial value calculation of the mirror evolution is likely to better capture the resulting dynamics than our model in which the anisotropy slowly increases in time.} Our results also show that the dominant wavelength is initially close to a few ion Larmor radii, with a subsequent migration to longer wavelengths through a coalescence process. This wavelength growth is consistent with observations \\citep[e.g.,][]{TsurutaniEtAl11b, SchmidEtAl14} and with theoretical models that predict the growth of mirror mode wavelengths due to magnetic diffusion \\citep{HasegawaEtAl11}.\\newline  \\noindent The fact that our $\\beta_i \\sim 1-6$ calculations find that the mirror instability dominates the nonlinear state of the fluctuations (even though the IC and mirror instabilities have comparable linear thresholds and growth rates for a bi-Maxwellian distribution) is also consistent with heliospheric observations. In particular, solar wind measurements at $\\beta_i\\sim 1-10$ show that $p_{\\perp,i}/p_{||,i}$ is limited by the linear threshold of the mirror instability \\citep[instead of the IC instability. See ][]{hellinger2006,bale2009}. The observations of \\citet{bale2009} also show enhanced magnetic fluctuations and magnetic compressibility that follow the linear mirror threshold fairly well. It is important to emphasize, however, that the observations presented by \\citet{hellinger2006} and \\citet{bale2009} showing dominance of the mirror instability also include the regime $\\beta_i \\ll 1$, which is not considered in our study. Also, the dominance of the mirror instability could be affected by properties of the solar wind that are usually not included in calculations of instability thresholds. These include the presence of $\\sim 5\\%$ of alpha particles \\citep[see, e.g., Figure 7 of ][]{MatteiniEtAl12}, the drift of alpha particles with respect to protons \\citep{MarschEtAl82a}, and the frequent presence of a proton beam \\citep{MarschEtAl82b}.\\newline  \\noindent We interpret the dominance of the mirror modes in our simulations as being due to the suppression of the IC waves, which is caused by the ion velocity distribution function departing from bi-Maxwellian \\citep{Isenberg12,Isenberg13}. This interpretation also requires the IC modes to have finite amplitudes (although smaller than the one of the mirror modes) so that they can maintain, via pitch-angle scattering, an ion distribution function that is marginally stable to the growth of IC waves.  To the best of our knowledge, these IC modes have not been found by in-situ heliospheric measurements during mirror dominated events. This is consistent with the fact that in our simulations the mirror and IC modes only have comparable amplitudes during the initial, exponential regime (Fig. \\ref{fig:beta6}$a$ and $b$). In the saturated regime (which is likely to be more representative of the observed state of the plasma in the heliospheric and astrophysical systems), we find that the mirror modes have significantly larger amplitudes. Moreover, this dominance increases with the magnetization  $\\omega_{c,i}/s$, as can be seen by comparing Figures \\ref{fig:beta6}$a$ and $b$. \\footnote{This dependence on $\\omega_{c,i}/s$ can be accounted for by the fact that the scattering rate of cyclotron-frequency fluctuations is expected to be $\\sim \\omega_{c,i} (\\delta B/B)^2$. If we equate this to the shear parameter $s$ (which is, roughly, the rate at which pitch-angle scattering takes place, as can be seen from the changing rate of $\\mu$ in Fig. \\ref{fig:muderivs}), one obtains that $(\\delta B/B)^2 \\sim s/\\omega_{c,i}$.} Thus, for realistic values of $\\omega_{c,i}/s$, we do not anticipate easily detectable IC waves by in-situ measurements in cases where the mirror modes dominate. The departure from the bi-Maxwellian ion distribution function has not been confirmed by in-situ heliospheric measurements of mirror-dominated events either. We believe that this is due to the high sensitivity of resonant instabilities to gradients in the velocity distribution function. Thus, in cases where IC and mirror modes can have similar growth rates, only small deviations from the bi-Maxwellian distribution can significantly reduce the importance of the IC modes and give way to the dominance of the mirror modes. We, therefore, do not anticipate easily detectable departures from bi-Maxwellian associated to this effect. \\footnote{We notice that Helios measurements show distribution functions with significant deviations from the bi-Maxwellian shape in the fast solar wind, consistent with predictions based on the presence of IC waves \\citep{TuEtAl02, HeuerEtAl07}. Although these measurements do not correspond to the small deviations needed to suppress the IC modes that we consider here, they do show the strong effect that high amplitude IC modes can have on the ion distribution.} Nevertheless, future careful observations comparing the nature of the fluctuations and the ion distribution functions in the solar wind can help to unveil these effects.\\newline  \\noindent A second application of our results is to low-collisionality accretion flows onto black holes. In these systems pressure anisotropies can produce an anisotropic stress that can increase the transport of angular momentum and the heating of the plasma \\citep{SharmaEtAl06}. Our results are consistent with \\citet{SharmaEtAl06}'s hypothesis that the mirror and/or IC instabilities would provide enhanced pitch angle scattering and regulate the evolution of the plasma pressure tensor.  Axisymmetric kinetic simulations have also demonstrated this directly \\citep{RiquelmeEtAl12} though the restriction to axisymmetry limits the quantitative applicability of these initial simulations to real disks.  One of the key conclusions of our study in this paper is that even modest plasma magnetization is sufficient to quantitatively capture the nonlinear saturation of the mirror instability.  For example, our $\\beta_i = 20$ results in Figures \\ref{fig:beta20B2s} \\& \\ref{fig:mubeta20} show that the evolution of the ion pressure anisotropy, ion magnetic moment, and turbulent fluctuations are quite similar for $\\omega_{c,i}/s = 93$ and $\\omega_{c,i}/s = 670$.  The primary difference is that in the lower magnetization run the plasma has a larger ``overshoot'' of the saturation pressure anisotropy set by the linear mirror instability threshold.  This weak dependence on magnetization is encouraging because it is computationally infeasible to reach $\\omega_{c,i}/s = 670$ in three-dimensional kinetic simulations of astrophysical plasmas (e.g., accretion disks).\\newline  \\noindent One of the interesting questions not fully addressed by our work is the saturation of velocity space instabilities in turbulent low-collisionality plasmas where the shear rate of the mean magnetic field itself fluctuates in time and the magnitude of the large-scale turbulent fluctuations satisfies $\\delta B \\ll \\langle B \\rangle$.  Our calculations in this paper have focused on the case where an imposed velocity shear amplifies a background magnetic field by order unity or more.  If, on the other hand, the shear acts for less time such that the background field is amplified by significantly less, the plasma may never become mirror/IC unstable (depending on $\\beta_i$ and $\\delta B/B$) or the background shear may reverse sign before the mirror/IC fluctuations have reached sufficiently large amplitudes for pitch angle scattering to set in.  This regime will be explored in future work. \\newline   \\noindent Another key question not addressed by our study is isotropization of the electron pressure anisotropy. In future work we will study how fast electrons interact with the large-scale mirrors generated by the ion pressure anisotropy, and whether there are separate electron-scale fluctuations that produce significant electron pitch-angle scattering (e.g., the whistler instability). \\newline"
    },
    "1402/1402.0478_arXiv.txt": {
        "abstract": "We suggest a new Bayesian analysis using disjunct M-R constraints for extracting probability measures for cold, dense matter equations of state. One of the key issues of such an analysis is the question of a deconfinement transition in compact stars and whether it proceeds as a crossover or rather as a first order transition. The latter question is relevant for the possible existence of a critical endpoint in the QCD phase diagram under scrutiny in present and upcoming heavy-ion collision experiments. ",
        "introduction": "One of the most challenging problems of modern physics concerns the structure of the phase diagram of quantum chromodynamics (QCD). Experimental programs with heavy-ion collisions (HIC) at ultrarelativistic energies and large-scale simulations of lattice QCD at finite temperature are performed to identify the position and the character of the suspected transition from a gas of hadronic resonances to a quark-gluon plasma in these systems characterized by almost perfect symmetry between particles and antiparticles, i.e. vanishing baryon density. It is nowadays established that there is a crossover transition with a pseudocritical temperature of $T_c=(154 \\pm 9)$ MeV \\cite{Bazavov:2011nk} from lattice QCD and a chemical freeze-out temperature $T_{\\rm fo}=(156 \\pm 5)$ MeV \\cite{Becattini:2012xb} for all hadron species from a statistical model analysis of hadron production data at $\\sqrt{s}=2.7$ TeV from the ALICE experiment at CERN LHC.  Both, HIC experiments and ab-initio lattice QCD cannot address, however, the QCD phase diagram at low and vanishing temperatures where the QCD phase transition is eventually of first order so that a critical end point (CEP) of first order transitions would result. The position of such a CEP in the phase diagram would be a landmark for the studies of strongly interacting matter under extreme conditions as it could help to identify the universality class of QCD. However, the beam energy scan programs at RHIC (STAR experiment) or at CERN SPS (NA61 experiment) have not yet brought any evidence for the existence of the CEP. Moreover, some theoretical studies argue that it could be absent at all because of the persistence of repulsive vector meson mean fields in dense matter \\cite{Bratovic:2012qs}, see also the discussion in \\cite{Contrera:2012wj}.  In this situation, where effective model approaches give contradicting results, and new experimental facilities like NICA at JINR Dubna and the CBM experiment at FAIR Darmstadt are not yet operative, a guidance for progress in the field could come from astrophysics of compact objects, namely from the precise mass and radius measurement of pulsars  \\cite{Miller:2013tca}. The main question to be answered is \\cite{Blaschke:2013ana,Alvarez-Castillo:2013cxa}: Can there be a trace of a (strong) first order phase transition in cold nuclear (neutron star) matter in mass-radius data from compact star observations?  In the present contribution, we outline a model study for this case which is based on Bayesian analysis (BA) methods and how they could guide future pulsar observational campaigns. ",
        "conclusions": "As a result of the BA for the three-parameter EoS (\\ref{eq:hybrid}) each of the $10 \\times 10 \\times 10 = 1000$ points in the three-dimensional parameter space has been assigned a probability value according to (\\ref{p_event}). We have selected the six most likely parameter sets and show the corresponding equations of state in the upper right panel of Fig.~\\ref{AHP_Scheme}. The corresponding $M-R$ and $M-M_B$ plots are given in the lower right panels of Fig.~\\ref{AHP_Scheme}, together with the corresponding observational constraints.  Alternativey, one can present the probabilities as histograms in the parameter space. Here, we are restricted to a two-dimensional subset (LEGO plots). Therefore we choose the three classes of speed of sound values for which we found the maximum probabilities and depict for each of them the LEGO plot of probabilities in the subspace of critical energy density ($\\epsilon_c$) and strength of the first-order transiition ($\\gamma$) in Fig.~\\ref{Bayesian}. \\begin{figure}[ht!] \\includegraphics[width=0.9\\textwidth]{3LegoPlots.eps} \\caption{LEGO plots for Bayesian analysis in the plane of two parameters $\\left(\\epsilon_c,\\gamma\\right)$ for three fixed values of the third parameter ${c}_{s}^{2}$.} \\label{Bayesian} \\end{figure} It is interesting to note that for the stiffest high-density EoS ($c^2_{\\rm QM}=0.922$), we obtain a bimodality of the probability distribution. A crossover behaviour ($\\gamma=0$) has the same probability as a strong first order transition ($\\gamma=0.5$).  We conclude that the Bayesian Analysis we have presented here presents an alternative to the previously developed one in Ref.~\\cite{Steiner:2010fz}. It is based on three statistically independent constraints which are extremely selective and result in maximum probabilities of about 1\\% for the most likely EoS. Thus this method will have a sufficient selective power when applied to a broader class of EoS. This concerns in particular EoS being stiffer (with larger radii) on the hadronic side and with microscopically founded EoS on the high density (quark matter) side. To test the interesting possibility that a strong first order phase transition might occur in massive neutron stars of $2~M_\\odot$ \\cite{Antoniadis:2013pzd,Demorest:2010bx} and lead to the appearance of a ``third family'' (mass twins) of hybrid stars at this high mass, one should perform radius measurements for these massive neutron stars. If it might turn out that the corresponding radii might be significantly different, e.g., by $1-2$ km with a standard deviation of about 500 m, then one would be able to disselect EoS without a phase transition, given the very narrow range of uncertainty in the mass of these objects. This possibility offers bright prospects for future observational campaigns and bears the chance to ``prove'' the existence of a critical point \\cite{Blaschke:2013ana,Alvarez-Castillo:2013cxa} in the QCD phase diagram from astrophysical observations!"
    },
    "1402/1402.2682_arXiv.txt": {
        "abstract": "We present Submillimeter Array observations, in the 1.3\\,mm continuum and the \\chtcnte~ line of 17 hot molecular cores associated with young high-mass stars. The angular resolution of the observations ranges from 1\\farcs0 to 4\\farcs0. The continuum observations reveal large ($>$3500 AU) dusty structures with gas masses from 7 to 375 M$_\\odot$, that probably surround multiple young stars. The \\chtcn~line emission is detected toward all the molecular cores at least up to the $K=6$-component and is mostly associated with the emission peaks of the dusty objects. We used the multiple $K$-components of the \\chtcn~ and both the rotational diagram method and a simultaneous synthetic LTE model with the XCLASS program to estimate the temperatures and column densities of the cores. For all sources, we obtained reasonable fits from XCLASS by using a model that combines two components: an extended and warm envelope, and a compact hot core of molecular gas, suggesting internal heating by recently formed massive stars. The rotational temperatures lie in the range of 40-132\\,K and 122-485\\,K for the extended and compact components, respectively. From the continuum and \\chtcn~results, we infer fractional abundances from $10^{-9}$ to $10^{-7}$ toward the compact inner components, that increase with the rotational temperature. Our results agree with a chemical scenario in which the \\chtcn~molecule is efficiently formed in the gas phase above 100-300\\,K, and its abundance increases with temperature. ",
        "introduction": "Massive stars ($M>8$\\,\\mo) are born inside of dense cores located in large and massive molecular clouds \\citep[e.g.,][]{1999PASP..111.1049Garay,2005IAUS..227...59C}. These massive star-forming regions (MSFRs) have a substantial impact on the evolution of the interstellar medium (ISM) and make important contributions to its dynamics and chemistry. For example, molecular outflows, jets, stellar winds and supernovae associated with MSFRs push into their surroundings, promoting additional star formation and mixing the ISM.  One of the first manifestations of massive star formation is the so-called hot molecular core phase \\citep[HMCs;][]{2000prpl.conf..299Kurtz, 2005IAUS..227...59C}. This phase is characterized by molecular gas condensations at relatively high temperatures ($>$100\\,K) and high densities ($\\sim$10$^5$--10$^8$\\,\\cmth), associated with a compact ($<$ 0.1\\,pc), luminous ($>$~10$^4$\\,\\lo), and massive ($\\sim$10-1000\\,\\mo) molecular core.  HMCs show a forest of molecular lines, especially from organic species \\citep[e.g.,][]{2005ApJS..156..127C}. Many of these molecules probably were formed on grain mantles during a previous cold phase, while others were produced by gas-phase reactions after \\ca parents species\\cc~were evaporated from the grains by the strong radiation of embedded or nearby protostars \\citep[see ][]{2009ARA&A..47..427H}.  Both models and observations suggest that massive HMCs are collapsing and accreting mass onto a central source(s) at rates of $10^{-4}-10^{-3}$\\,\\mo\\,yr$^{-1}$ \\citep{2009ApJ...694...29O,2009ApJ...698.1422Z}. These intense mass accretion rates are high enough to prevent the development of an ionized region around the massive star(s) at least in the early stages \\citep{1999ApJ...525..808O}. Thus, HMCs probably precede ultracompact \\hii~regions \\citep[\\uchii;][]{2000prpl.conf..299Kurtz, 2001ApJ...550L..81Wilner}. Indeed, sub-arcsecond observations argue in favor of this scenario, particularly those showing embedded \\uchii-regions, strong (sub)millimeter emission from dust condensations, or strong mid-IR emission from internal objects \\citep[e.g.;][]{2010A&A...509A..50C,2011A&A...533A..73C}.  In the above scenario a HMC corresponds to the most internal clump of molecular material collapsing and probably feeding other structures and the massive stars inside \\citep{2005IAUS..227...59C,2001ApJ...550L..81Wilner}. However, recent sensitive high angular resolution observations suggest that the prototypical HMC, Orion BN/KL, may not follow this model. In this case, a close dynamical interaction of three young protostars produced an explosive flow and illuminated a pre-existing dense clump, thus creating the HMC \\citep{2011A&A...529A..24Z,2011ApJ...739L..13G}. Also, toward G34.26+0.15 (another prototypical HMC), \\citet{2007ApJ...659..447M} failed to find any embedded protostars within the hot core. The different nature of internally and externally heated HMC makes it important to distinguish between them.  With this in mind, we present a study using 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) archival observations of \\chtcnte~and 1.3\\,mm continuum emission, toward 17 MSFRs in the HMC stage. \\chtcn~(methyl cyanide) is frequently used as an effective thermometer and to estimate gas density toward HMCs \\citep[e.g.,][]{2005ApJS..157..279Araya, 2001ApJ...558..194P}. Our main goal is to use the same molecular tracer toward a relatively large group of sources to study the inner-most and hottest material, estimating densities, temperatures, masses, abundances, and the spatial distribution of the dust emission and \\chtcn~molecular gas.  In Section \\ref{observ} we describe the archival observations presented in this study. In Section \\ref{results} we report the results and analysis of the millimeter continuum data and the molecular line emission. In Section \\ref{comments} we comment briefly on each source, giving the physical characteristics from the literature and from our results. In Section \\ref{discussion} we discuss our results, first comparing the spatial distribution of the continuum emission and molecular emission, and then estimating the temperatures and densities of the regions from an LTE analysis of the \\chtcnte~ spectra and through the rotation diagram method. Finally, in Section \\ref{conclus}, we present our main conclusions. ",
        "conclusions": "We studied 17 hot molecular cores in the \\chtcnte~lines and the 1.3\\,mm continuum. The sources were observed with the SMA at $\\sim$220\\,GHz, with either the compact or extended configuration.  From the 1.3\\,mm continuum, we detected dusty structures with physical sizes of 0.01--0.1\\,pc, gas masses of 7--375\\,\\mo, and column densities of 0.1--6.7\\x10$^{24}$\\,\\cmtw. The continuum emission coming from dust ranges from 8\\% to 98\\% of the total flux.  All 17 sources show multiple molecular lines but different molecular richness. All sources show five or more $K$-components of \\chtcnte. Some spectra showed emission up to the $K=8$-component, which traces gas at $\\sim$525\\,K.  Based on these emission lines we estimated rotational temperatures, column densities, and fractional abundances, using both rotation diagrams and the XCLASS program that generates synthetic spectra. From the rotation diagram method we find temperatures from 90 to 500\\,K, and column densities from 2.5\\x10$^{13}$ to 2.5\\x10$^{16}$\\,\\cmtw. With XCLASS we find temperatures from 40 to 132\\,K and column densities from 1.6\\x10$^{13}$ to 1.0\\x10$^{15}$\\,\\cmtw~for the extended component, and temperatures from 122 to 485\\,K and column densities from 6.8\\x10$^{14}$ to 5.1\\x10$^{17}$\\,\\cmtw~for the compact component.  We used the rotation temperatures estimated with XCLASS to derive the gas mass from the 1.3\\,mm continuum. With the multiple $K-$lines of \\chtcn~we find a good fit between observed and synthetic spectra for the two-component XCLASS model; this result, and a close spatial coincidence between the molecular gas and the continuum emission, suggest that most of these HMC are internally heated. Sub-arcsecond observations are necessary to explore their structure in greater detail.  The fractional abundance of \\chtcn~toward the hot-inner components shows a marked increase with temperature. This can be understood if we consider that \\chtcn~molecules form in the hot gas phase when parent N-bearing species, such as \\nht, are evaporated from grain mantles."
    },
    "1402/1402.0687_arXiv.txt": {
        "abstract": "We present the results of a Keck/DEIMOS survey of Ultra Compact Dwarfs (UCDs) in the Perseus Cluster core. We confirm cluster membership for 14 UCDs, with radial velocities $\\sim5300$~km~s$^{-1}$. Two of these confirmed Perseus UCDs have extremely blue colours ($B-R < 0.6$~mag), reside in star forming filaments surrounding NGC~1275, and have likely formed as massive star clusters in the last $\\sim100$~Myr. We also measure a central velocity dispersion of a third, UCD13 ($\\sigma_{0} = 38 \\pm 8$~km~s$^{-1}$), the most extended UCD in our sample. We determine it to have radius $R_{e} = 85 \\pm 1.1$~pc, a dynamical mass of $(2.3 \\pm 0.8)\\times10^{8}$~M$_{\\sun}$, and a metallicity [Z/H]$ = -0.52^{+0.33}_{-0.29}$~dex. UCD13 and the cluster's central galaxy, NGC 1275, have a projected separation of 30 kpc and a radial velocity difference of $\\sim20$~km~s$^{-1}$. Based on its size, red colour, internal velocity dispersion, dynamical mass,  metallicity and proximity to NGC~1275, we argue that UCD13 is  likely the remnant nucleus of a tidally stripped dE, with this progenitor dE having $M_{B} \\approx -16$~mag and mass $\\sim10^{9}$~M$_{\\odot}$. ",
        "introduction": "Ultra Compact Dwarfs (UCDs) were first identified over a decade ago (Hilker et al. 1999; Drinkwater et al. 2000), but whether they should be considered as galaxies or massive star clusters is still uncertain (e.g. Hasegan et al. 2005; Mieske et al. 2009; Hilker 2009; Forbes \\& Kroupa 2011; Willman \\& Strader 2012). Brodie et al. (2011) adopted a working definition for UCDs to have half-light sizes of $10~{\\rm pc} < R_{h}<100$~pc and $M_{V} < -8$~mag (with no upper limit on luminosity).  The two main formation channels discussed in the literature are variations of merged star clusters \\citep{1998MNRAS.300..200K,2002MNRAS.330..642F,2011A&A...529A.138B} or stripped dwarf galaxies \\citep{1994ApJ...431..634B,2001ApJ...552L.105B}. In this sense their current size and luminosity are thought to be the result of a `bottom-up' growth from lower masses, or a `top-down' evolution from higher masses. Support for the idea of both formation channels for UCDs comes from the results of  Norris \\& Kannappan (2011). They found many UCDs to be consistent with high mass globular clusters. However, they also reported one UCD around NGC 4546 that was in a counter-rotating orbit with respect to the host galaxy and its globular cluster system. Furthermore, they calculated that the UCD had a very short dynamical friction timescale. This strongly suggested that it had been recently captured and was not part of the galaxy's original GC system.  Evidence for UCDs forming as massive star clusters in the nearby Universe was presented by Penny et al. (2012), who found young, blue candidate UCDs with ($B-R$)$_{0}<0.6$~mag in the Perseus Cluster associated with gas filaments around the cluster's central galaxy NGC 1275. NGC1275 therefore provides a location in which to study the formation of massive star clusters and proto-UCDs in the nearby Universe. If these massive blue star clusters are able to remain bound against the cluster tidal potential, they will be seen as typical red UCDs once their stellar populations fade in a timescale of $\\sim500$~Myr.  Using a state-of-the-art compilation of pressure-supported systems with old stellar populations, Brodie et al. (2011) found that UCDs had a tight colour-magnitude relation that was different from normal globular clusters. This, and kinematic differences, argued against the merged star cluster scenario. However, they also found that UCDs do not have a well-defined size-luminosity distribution as might have been expected in the stripping scenario \\citep{2013MNRAS.433.1997P}. A further expectation of the stripping scenario is that the remnants come from a range of progenitor masses. Hence we would expect a continuity of objects between the most massive UCDs and compact ellipticals (cEs), the latter of which are thought to be the tidally stripped remnant of a normal galaxy (Faber 1973). Such intermediate mass objects are too massive to be explained as merged star clusters.   In the compilation of Brodie et al. (2011), their figure 8 shows very few objects with confirmed sizes and luminosities intermediate between those of UCDs and cEs. One such object is M59cO highlighted by Chilingarian \\& Mamon (2008) as a ``missing link'' between UCDs and cEs. This object, with a confirmed velocity placing it in the Virgo cluster, has  $M_{B} = -12.34$~mag ($M_V$ $\\sim -13.0$~mag). A dual Sersic fit to the surface brightness profile gave a `nucleus' of half-light radius $R_{e}=13$~pc with an extended component of $R_{e}=50$~pc. The object was also found to be old and relatively metal-rich. Another extreme UCD in the Virgo Cluster, M60-UCD1 \\citep{2013ApJ...775L...6S} with $M_{B} = -13.24$~mag and a dynamical mass $M_{dyn} = (2.0 \\pm 0.3) \\times 10^{8}$~M$_{\\sun}$ is the densest galaxy identified to date, and occupies the region of size-magnitude space between the most massive globular clusters and the faintest compact elliptical galaxies.  In this work we present the results of a Keck DEIMOS follow-up study of UCD candidates in the nearby Perseus Cluster ($D=71$~Mpc) identified in \\citet{Penny12} out to 250~kpc from the cluster centre. The Brightest Cluster Galaxy (BCG) is the unusual NGC 1275, which is actively accreting/disrupting satellite galaxies at the present time, making Perseus an ideal environment in which to study the formation of UCDs as disrupted dEs vs. massive star clusters. The focus on this paper is the origin of UCDs in clusters- are they massive star clusters or the cores of tidally stripped galaxies? In particular we examine on the origin of the largest Perseus UCD, UCD13, via its size, kinematics, cluster location and stellar population.  This paper is organised as follows. In Section~\\ref{sec:data}, we describe our observations and their reduction. In Section~\\ref{perucd} we briefly discuss the properties of the confirmed Perseus Cluster members, including the size luminosity distribution for UCDs. An analysis of UCD13 is presented in Section~\\ref{ucd13}, with a tidal stripping formation scenario for UCD13 examined in Section~\\ref{sec:ucd13}. A star cluster origin for UCD13 is discussed in Section~\\ref{starcluster}. We discuss our results in Section~\\ref{sec:discuss}, and present our conclusions in Section~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  We have examined the properties of Perseus UCDs with their cluster membership confirmed via \\textsc{keck deimos} spectroscopy. Perseus UCDs with sizes $<20$~pc are consistent with being massive globular clusters or the nuclei of stripped dwarf ellipticals. Two blue UCDs with colours ($B-R$)~$<0.6$ are confirmed as members of the Perseus Cluster, and must have formed within the last 100~Myr as massive star clusters in the star forming filaments surrounding NGC 1275.  The largest UCD (UCD13) is likely the nuclear remnant of a tidally stripped dwarf elliptical that has retained a large fraction of its stellar envelope, consistent with the simulations of \\citet{2013MNRAS.433.1997P}. Its dynamical mass is $(2.3 \\pm 0.8)\\times10^{8}$~M$_{\\sun}$, making it one of the most massive UCDs identified to date. Its internal velocity dispersion, $\\sigma_{0} = 38 \\pm 9$~km~s$^{-1}$, along with its metallicity $\\rm{[Z/H]} = -0.52^{+0.33}_{-0.29}$~dex, and colour $(B-R)_{0} = 1.37$, are consistent with its origin as a dE with an $M_{B}$ of $-16$ and a mass $\\sim10^{9}$~M$_{\\odot}$. Its size ($R_{e} = 85 \\pm 1.1$~pc) and location at the centre of Perseus, the UCD is inconsistent with being a massive star cluster."
    },
    "1402/1402.6965_arXiv.txt": {
        "abstract": "We propose a scenario where the saxion dominates the energy density of the Universe and reheats the standard model sector via the dilatonic coupling, while its axionic partner contributes to dark matter decaying into photons via the same operator in supersymmetry. Interestingly, for the axion mass $m_a \\simeq 7$\\,keV and the decay constant $f_a \\simeq 10^{14-15}$\\,GeV, the recently discovered X-ray line at $3.5$\\,keV in the XMM Newton X-ray observatory data can be explained. We discuss various  cosmological aspects of the $7$\\,keV axion dark matter such as the production of axion dark matter, the saxion decay process, hot dark matter and isocurvature constraints on the axion dark matter, and the possible baryogenesis scenarios. ",
        "introduction": "In supergravity and superstring theories there appear many moduli fields at low energy scale through compactifications of extra dimensions~\\cite{fluxc}. Moduli fields must be stabilized to obtain a sensible low-energy theory, and it is known that many of them are fixed by flux compactifications and acquire a heavy mass~\\cite{Grana:2005jc}. The remaining light moduli not fixed by the fluxes can be stabilized either by instantons/gaugino condensations {\\it a la} KKLT~\\cite{Kachru:2003aw} or by supersymmetry (SUSY) breaking effects~\\cite{Berg:2005yu,Balasubramanian:2005zx,Conlon:2006tq,Choi:2006za}.\\footnote{ For instance the QCD axion could be the axion component of such a modulus field mainly stabilized by the SUSY breaking effects.} Such  light moduli fields  may play an important role in  cosmology; some of them may dominate the Universe and decay into the standard model (SM) sector, or others could contribute to dark matter or dark energy if their masses are sufficiently light.     Recently an unidentified X-ray line at about $3.5$~keV in the XMM-Newton X-ray observatory data of various galaxy clusters and the Andromeda galaxy was reported independently by two groups~\\cite{Bulbul:2014sua,Boyarsky:2014jta}. While there are a variety of systematic uncertainties that can affect the observed line energy and flux, it is intriguing that the X-ray line can be explained by decaying dark matter such as sterile neutrinos\\footnote{ Recently, Ishida and two of the present authors (KSJ and FT) showed that the small mass and mixing of sterile neutrino dark matter suggested by the X-ray line can be easily realized by the split flavor mechanism where the breaking of flavor symmetry is tied to the breaking of the $B-L$ symmetry~\\cite{Ishida:2014dlp}. }~\\cite{Dolgov:2000ew} or moduli fields~\\cite{Hashiba:1997rp,Asaka:1997rv,Kasuya:2001tp,Kusenko:2012ch}.  The observations suggest the mass and the lifetime of the dark matter as~\\cite{Bulbul:2014sua,Boyarsky:2014jta}: \\bea m_{\\rm DM} &\\simeq&7 {\\rm \\, keV},\\\\ \\tau_{\\rm DM} &\\simeq& 2 \\times 10^{27}  - 2 \\times 10^{28}\\, {\\rm sec}, \\label{tau-obs} \\eea where we have used the values obtained by the M31 data~\\cite{Boyarsky:2014jta}, and we adopt them as reference values in the following analysis assuming that decaying dark matter is the origin of the $3.5$\\,keV X-ray line.  The light dark matter mass about $7$\\,keV may be due to some approximate symmetry forbidding the mass. We focus on the axion component of a modulus field $\\Phi = (\\sigma + i a)/\\sqrt{2}$ stabilized by SUSY breaking effects, where $\\sigma$ and $a$ are the saxion and the axion components, respectively. The axion $a$ can remain extremely light as a result of the axionic shift symmetry, \\beq \\Phi \\rightarrow \\Phi +  i C, \\label{shiftsym} \\eeq where $C$ is a real transformation parameter. The axion can acquire a small but non-zero mass of $7$\\,keV from some non-perturbative effects which explicitly break the above shift symmetry. We shall see that, if the modulus field $\\Phi$ is coupled to the SM gauge fields with a decay constant of order $\\GEV{14-15}$, the lifetime of the axion falls in the  range of (\\ref{tau-obs}), explaining the observed X-ray line. On the other hand, the saxion $\\sigma$ generically acquires a mass of order of the gravitino mass from SUSY breaking effects. The gravitino mass is not known, but it must be heaver than the electroweak scale in the gravity or anomaly mediation. We assume this is the case throughout this letter.  The mass hierarchy  between the saxion and the axion leads to a unified picture of the cosmological role of light moduli fields:  the saxion dominates the Universe and reheats the SM sector via the dilatonic coupling, while the axion contributes to dark matter decaying into photons via the same operator in SUSY. As we shall see shortly, the right abundance of axion dark matter can be produced by coherent oscillations for the saxion mass about $\\GEV{6}$ and the decay constant $f_a \\simeq \\GEV{14-15}$ without fine-tuning of the initial misalignment angle.  We shall also see that the axions are generically produced by the saxion decay, which may contribute to hot dark matter (HDM) component in agreement with the recent observations~\\cite{Wyman:2013lza,Hamann:2013iba, Battye:2013xqa}. Therefore, the detailed study of the decaying axion dark matter via the X-ray observation and the observations of large-scale structure can be a probe of not only the nature of dark matter but also the reheating of the Universe as well as the high-energy physics close to the GUT scale.    In this letter we propose a scenario in which the $7$\\,keV axion dark matter decaying into photons explains the origin of the $3.5$\\,keV X-ray line, while the saxion dominates the Universe and reheats the SM sector via the same dilatonic coupling in SUSY. We will study various aspects of this scenario, focusing on the saxion cosmology, the production mechanism of the axion dark matter,  the isocurvature  and HDM constraints, and possible baryogenesis scenarios in turn.  Lastly let us briefly mention the differences of our work from Ref.~\\cite{Kusenko:2012ch}  (and other works~\\cite{Hashiba:1997rp,Asaka:1997rv,Kasuya:2001tp}). One of the main differences is the SUSY breaking scale, i.e., the gravitino mass.  They focused on the light gravitino mass between keV and MeV, and consider moduli dark matter with a similar mass, which corresponds to the real component of the moduli, i.e.  the saxion, in our scenario. On the other hand, it is its axionic partner that becomes dark matter in our scenario. As long as the implications for the observation of the X-ray line are concerned, there is no significant difference between these two models.   The crucial difference is that the heavy gravitino we consider enables a scenario in which the saxion dominates and reheats the Universe via the same dilatonic coupling in SUSY. Then we can unambiguously discuss the saxion and axion cosmology. ",
        "conclusions": "There appear many moduli fields in the low energy through compactifications of extra dimensions in superstring theories. Some of the moduli fields may remain light after the closed string flux is turned on. We have focused on a modulus field which is stabilized by the SUSY breaking effect and its axion component remains much lighter than the saxion component. As long as the strong CP problem is solved by the string-theoretic QCD axion, there must be at least one such modulus field, and  in general, there might be more. As such moduli fields tend to be lighter than those stabilized by the non-perturbative effects {\\it a la} KKLT, they likely play an important cosmological role.  We have proposed a scenario in which the saxion component of such modulus field dominates the energy density of the Universe and reheats the SM sector via its dilatonic couplings, while its axion partner contributes to dark matter decaying into photons via the same dilatonic coupling to photons in SUSY. The point is that both the reheating of the Universe and the decay of dark matter into photons are induced by the same supermultiplet (i.e. saxion and axion) through the same operator in SUSY. This observation partially explains why dark matter decays into photons at all. If there are light axions, one of them can easily explain the dark matter abundance as  the axions are copiously generated by coherent oscillations. Then, there is no special reason why the axion dark matter should be coupled to photons. The situation changes if the bosonic partner, the saxion, dominates the Universe and reheats the SM sector through the same operator in SUSY. In this case, the axion dark matter must be coupled to the SM sector, in order for successful reheating. In other words, the decaying dark matter can be a probe of the reheating of the Universe.   We have also discussed the saxion decay process, the HDM constraint on the axions produced by the saxion decay, the isocurvature constraint on the axions produced by coherent oscillations, and the baryogenesis scenarios. Some of our results, especially those about the nature of axion dark matter (i.e. abundance, lifetime and isocurvature constraints), can be straightforwardly applied to the case in which the saxion does not dominate the Universe. This is likely the case e.g. if the Hubble parameter during inflation is smaller than the saxion mass.     Interestingly, for the axion mass $m_a \\simeq 7$\\,keV and the decay constant $f_a \\simeq 10^{14-15}$\\,GeV,   the recently discovered X-ray line at $3.5$\\,keV in the XMM Newton X-ray observatory data  can be explained by the decay of the axion dark matter. The suggested value of the decay constant is within the expected range for the string-theoretic axion. It is of course possible to consider field-theoretic axions or pseudo Nambu-Goldstone bosons of mass $7$\\, keV  which have couplings to photons with a similar strength. The detailed  X-ray line search in future may not only probe the nature of dark matter but also unravel the very early history of our Universe as well as physics close to the GUT scale."
    },
    "1402/1402.3637_arXiv.txt": {
        "abstract": "We present deep VLA observations of the polarization of radio sources in the GOODS-N field at 1.4~GHz at resolutions of 1.6\\arcsec and 10\\arcsec. At 1.6\\arcsec, we find that the peak flux cumulative number count distribution is N($>$p) $\\sim$  45 * (p/30$\\mu$Jy)$^{-0.6}$  per square degree above a detection threshold of 14.5~$\\mu$Jy. This represents a break from the steeper slopes at higher flux densities, resulting in fewer sources predicted for future surveys with the SKA and its precursors. It provides a significant challenge for using background RMs to study clusters of galaxies or individual galaxies.  Most of the polarized sources are well above our detection limit, and are radio galaxies which are well-resolved even at 10\\arcsec, with redshifts from $\\sim$0.2 - 1.9.  We determined a total polarized flux for each source by integrating the 10\\arcsec~polarized intensity maps, as will be done by upcoming surveys such as POSSUM. These total polarized fluxes are a factor of 2 higher, on average, than the peak polarized flux at 1.6\\arcsec; this would increase the number counts by $\\sim$50\\% at a fixed flux level. The detected sources have rotation measures (RMs) with a characteristic rms scatter of $\\sim$11$\\frac{rad}{m^2}$ around the local Galactic value, after eliminating likely outliers. The median fractional polarization from all total intensity sources does not continue the trend of increasing at lower flux densities, as seen for stronger sources.  The changes in the polarization characteristics seen at these low fluxes likely represent the increasing dominance of star-forming galaxies. ",
        "introduction": "Deep surveys for polarized radio sources are motivated by several different scientific goals.   One is to use their rotation measures (RMs) as a background grid to study the magneto-ionic structure of the Milky Way \\citep{taylornvss}, or nearby galaxies \\citep{step08}, or clusters of galaxies \\citep{govoni10}, and in the unique case of Centaurus~A, even the lobes of radio galaxies \\citep{fornax}.   The most extensive RM grid currently available is that of \\cite{taylornvss}, with 37,543 polarized sources reconstructed from the NVSS \\citep{NVSS} survey.   It has a source density of $\\sim$1 per square degree, making it an excellent probe of the large scale structure of the Milky Way, but insufficient for studies of even nearby extended extragalactic objects.   This has motivated ambitious plans for polarization surveys such as POSSUM \\citep{possum1} on the Australian SKA precursor (ASKAP, \\cite{askap} ), with a target of polarized source densities $\\sim$100 times higher than the NVSS.  A second motivation for deep polarization surveys is to probe source populations, e.g., as a function of flux level, luminosity and redshift.  Polarization counts offer complementary information to counts in total intensity (e.g., \\cite{icounts,morrisonG}), and can help distinguish between the contributions of star-forming galaxies, which are expected to have low polarizations, from AGN radio sources.  Finally, polarization surveys can inform the physical models of radio galaxies, e.g., the amount of magnetic field order/disorder,  the presence of an external magnetised medium and even the interactions  of the thermal and relativistic plasmas \\citep{guidetti12}.  For these purposes, the RMs, the depolarization as a function of wavelength, and eventually the spatially unresolved Faraday structure, e.g. \\cite{law}, are important diagnostics.  The deepest polarization surveys to date are those of the ELAIS-N1 field by \\cite{elais2}, which reach down to polarized fluxes p$>$270~$\\mu$Jy, and the ATLBS field studied by \\cite{atlbs}.\\footnote{ As discussed below, we consider the latter work reliable only for p$>$1000~$\\mu$Jy.}  The very sensitive 1.4~GHz Karl G. Jansky Very Large Array (VLA) survey of the GOODS-N field \\citep{morrisonG}, with an rms sensitivity of several $\\mu$Jy, provided an excellent opportunity to extend polarization studies to much lower levels, albeit over a smaller area of the sky.   A parallel effort to reach polarized sensitivity levels in the $\\mu$Jy range is based on the stacking of sources from the NVSS \\citep{stack}, and is complementary to the work presented here.  This paper is arranged as follows.   Section II describes the observations and data reduction, while Section III describes the methodology behind our source detection and presents the results. More detailed information on individual sources is presented in Section IV.   In Section V, we derive the polarized number counts and fractional polarization distributions as a function of flux density, and compare these to results at higher fluxes and models seeking to predict the behavior at $\\mu$Jy levels.  Section VI contains a discussion of the results, in particular their implications for some key experiments planned for the next generation of radio telescopes. ",
        "conclusions": "Deep observations of the GOODS-N field at 1.4~GHz using the VLA were used to conduct an automated search for polarization from total intensity sources at 1.6\\arcsec~ resolution, and a visual search of extended sources at 10\\arcsec~ resolution.  $\\bullet$ For the automated search, where the completeness, contamination and available search area as a function of peak polarized flux level are well characterized, we detect a flattening of the cumulative polarized number counts down to 14.5~$\\mu$Jy. Above P=1~mJy, other work shows a cumulative slope $d$log(N($>$P))/$d$log(p)$\\propto$P$^{-1.4}$, while for P$<$1~mJy, we find $d$log(N($>$P))/$d$log(p)$\\propto$P$^{-0.6}$.  Most of the GOODS-N detections are from extended sources with $\\sim$0.2$<$z$<$1.9, and are much brighter than the survey detection threshold.  $\\bullet$  The predicted number of polarized sources for upcoming surveys such as POSSUM, with a 5$\\sigma$ threshold of 50~$\\mu$Jy, is $\\sim$35$\\pm$10 per square degree, based on the 1.6\\arcsec~ counts. Integrated polarized fluxes at 10\\arcsec~ resolution may be a factor of two higher, boosting the counts. The rms dispersion in RM, for most of the sources, is 11\\rpm~ around the Galactic foreground level;  a small fraction of the sources have high values of $|$RM$|$, making them unsuitable for study of foreground Faraday screens.  At these source densities, individual source foreground studies will only be possible for the nearest galaxies and clusters of galaxies.  $\\bullet$ We also find, for 0.5~mJy~$<$~I~$<$~5~mJy, an end of the trend seen at higher fluxes for the fractional polarization to rise with decreasing total intensity.  This, and the number counts, are consistent with the dropoff in the radio galaxy population and the increasing dominance of star-forming sources in the surveys.  For I$<$100$\\mu$Jy, we can establish only weak upper limits on the median fractional polarization.  $\\bullet$ The detailed distributions of polarized number counts and fractions are dependent on whether peak fluxes or some form of integrated polarized fluxes are used, and on the angular resolution of the survey;  lower resolution surveys can pick up lower surface brightness polarized emission, but can also suffer from depolarization from changes of polarization position angle (intrinsic or RM-caused) within the beam.  \\vskip 1in"
    },
    "1402/1402.6823_arXiv.txt": {
        "abstract": "The gravitational lensing effects in the weak gravitational field by exotic lenses have been investigated intensively to find nonluminous exotic objects. Gravitational lensing based on $1/r^n$ fall-off metric, as a one-parameter model that can treat by hand both the Schwarzschild lens (n=1) and the Ellis wormhole (n=2) in the weak field, has been recently studied. Only for $n=1$ case, however, it has been explicitly shown that effects of relativistic lens images by the strong field on the light curve can be neglected. We discuss whether relativistic images by the strong field can be neglected for $n>1$ in the Tangherlini spacetime which is one of the simplest models for our purpose. We calculate the divergent part of the deflection angle for arbitrary $n$ and the regular part for $n=1$, $2$ and $4$ in the strong field limit, the deflection angle for arbitrary $n$ under the weak gravitational approximation. We also compare the radius of the Einstein ring with the radii of the relativistic Einstein rings for arbitrary $n$. We conclude that the images in the strong gravitational field have little effect on the total light curve and that the time-symmetric demagnification parts in the light curve will appear even after taking account of the images in the strong gravitational field for $n>1$. ",
        "introduction": "Gravitational lensing is useful to survey nonluminous objects for example extrasolar planets and dark matters (see~\\cite{Schneider_Ehlers_Falco_1992,Perlick_2004_Living_Rev,Perlick_2010,Petters_Levine_Wambsganss_2001, Schneider_Kochanek_Wambsganss_2006,Wambsganss_1998,Bartelmann_2010} for the details of the gravitational lensing and references therein). The gravitational lensing effects of the Schwarzschild lens in the weak gravitational field have been investigated for a hundred years.  Exotic objects such as wormholes, cosmic strings also cause gravitational lensing effects. The gravitational lenses in wormhole spacetimes were pioneered by Kim and Cho~\\cite{Kim_Cho_1994} and Cramer~\\textit{et al.}~\\cite{Cramer_Forward_Morris_Visser_Benford_Landis_1995} and then the gravitational lensing effects in various wormhole spacetimes have been investigated~\\cite{Safonova_Torres_Romero_2001_Jan, Eiroa_Romero_Torres_2001,Safonova_Torres_Romero_2001_Nov,Safonova_Torres_2002,Nandi_Zhang_Zakharov_2006,Rahaman_Kalam_Chakraborty_2007, Dey_Sen_2008,Takahashi_Asada_2013,Kuhfittig:2013hva} (See Visser~\\cite{Visser_1995} for the detail of wormholes).  The Ellis wormhole~\\cite{Ellis_1973,Bronnikov_1973} is the simplest and earliest wormhole of the Morris-Thorne class~\\cite{Morris_Thorne_1988,Morris_Thorne_Yurtsever_1988}. Recently, Bronnikov~\\textit{et al.} have shown that the Ellis wormhole metric supported by a perfect fluid with negative density and a source-free radial electric or magnetic field is linearly stable under the spherically symmetric perturbations and axial perturbations~\\cite{Bronnikov:2013coa} while the Ellis wormhole with the same metric supported by a phantom scalar field is unstable~\\cite{Armendariz-Picon_2002, Shinkai_Hayward_2002,Gonzalez_Guzman_Sarbach_2008_I,Gonzalez_Guzman_Sarbach_2008_II, Doroshkevich_Hansen_Novikov_Shatskiy_2009,Bronnikov:2011if}. The former is the first example of a stable wormhole which is not supported by thin shells in general relativity. The gravitational lensing in the Ellis wormhole spacetime has been investigated intensively~\\cite{Chetouani_Clement_1984,Perlick_2004_Phys_Rev_D, Nandi_Zhang_Zakharov_2006,Dey_Sen_2008,Abe_2010,Bhattacharya_Potapov_2010,Toki_Kitamura_Asada_Abe_2011,Nakajima_Asada_2012,Gibbons_Vyska_2012, Tejeiro_Larranaga_2012,Tsukamoto_Harada_Yajima_2012,Tsukamoto_Harada_2013,Yoo_Harada_Tsukamoto_2013,Takahashi_Asada_2013} because of its simplicity and its features which are caused by the gravitational potential which decreases as the inverse square of a radial coordinate $r$ under the weak-field approximation. Abe found numerically that the light curves of the microlensing by the Ellis wormhole show the time-symmetric demagnification under the weak-field approximation~\\cite{Abe_2010}. This means that the theorem that total magnification is always larger than the unity is true for mass lenses with the gravitational potential which has the form of $1/r$ but it is not true for exotic lens objects with the gravitational potential which has the form of $1/r^{2}$ under the weak-field approximation.  Gravitational lensing based on $1/r^n$ fall-off metric, as a one-parameter model that can treat by hand both the Schwarzschild lens $(n=1)$ and the Ellis wormhole $(n=2)$ in the weak field, has been recently studied~\\cite{Tsukamoto_Harada_2013,Kitamura_Nakajima_Asada_2013, Izumi_Hagiwara_Nakajima_Kitamura_Asada_2013,Kitamura_Izumi_Nakajima_Hagiwara_Asada_2013,Nakajima:2014nba}. In particular, Kitamura~\\textit{et al.} showed that the demagnification of the light curves appears in the case $n>1$~\\cite{Kitamura_Nakajima_Asada_2013}. This implies that the surveys of microlensing events are a good way to search exotic objects including the Ellis wormhole in our Galaxy and so on. However, these studies concentrate on the weak gravitational field and do not cover the gravitational lensing effects in the strong gravitational field.  The timelike and null geodesics in the Schwarzschild spacetime were studied by Hagihara~\\cite{Hagihara_1931} and Darwin~\\cite{Darwin_1959,Darwin_1961}. Darwin pointed out the existence of the relativistic images which are a series of faint images lying just outside the photon sphere~\\cite{Darwin_1959} in the Schwarzschild spacetime. A countably infinite number of relativistic images are formed in a spherically symmetric static spacetime, if the radius of the lens object is smaller than a photon sphere ($=3r_{g}/2$, where $r_{g}$ is the radius of the event horizon in the Schwarzschild spacetime)~\\cite{Bozza_2002,Hasse_Perlick_2002,Perlick_2004_Phys_Rev_D}. The gravitational lensing in the strong gravitational field by various black holes and wormholes has been investigated intensively in the recent decade (see~\\cite{Virbhadra_Ellis_2000,Virbhadra_Keeton_2008,Virbhadra_2009,Bozza_2010,Bozza_Mancini_2012,AzregAinou:2012xv} and references therein). In a series of the papers on the relativistic images, the observables for the suppermassive object at the center of our Galaxy with instruments which can measure the relativistic images are discussed intensively.  The Tangherlini solution is an exact solution of the Einstein equation in every dimension, where the solution is spherically symmetric, static and asymptotically flat. Here, as a generalization of the Einstein-Hilbert action to every dimension, the action for gravity is assumed to be the Ricci scalar even in $d \\neq 4$ dimensions for its simplicity, though the physics behind this assumption in higher dimensions is not clear. Then, the variation of the gravity action with respect to the metric provides the vacuum field equation as \"Ricci = 0\". It follows that the Tangherlini solution recovers the Schwarzschild solution in four dimensions.  An exact form of the metric with $1/r^n$ in the weak field limit would enable us to investigate such exotic lenses not only in the weak field but also in the strong field. Only for $n=1$ case, it has been explicitly shown that effects of relativistic lens images by the strong field on light curves can be neglected~\\cite{Bozza_Capozziello_Iovane_Scarpetta_2001}. For $n > 1$, this issue has not been addressed yet. Therefore, the main purpose of the present paper is to discuss whether relativistic images by the strong field can be neglected for $n > 1$. It is likely that a spacetime geometry with $1/r^n$ fall-off in the weak field is not unique but has many variants. The Tangherlini solution is one of the simplest models for our purpose. For recent investigations of exotic gravitational lenses, therefore, it would be of physical interest to examine the Tangherlini lens both in the weak field and in the strong one.  The gravitational lens in the strong field limit is related to the other phenomena such as the quasinormal modes of a black hole \\cite{Stefanov:2010xz,Wei:2013mda} and the high-energy absorption cross section \\cite{Wei:2011zw} which are caused by the nature of the null geodesic near the photon sphere. Thus, the investigation of gravitational lensing effects of the all-dimensional black hole in the strong field limit would give us a new perspective on the all-dimensional black hole.  This paper is organized as follows. In Sec.~II, we review the null geodesic in the Tangherlini spacetime and investigate the deflection angle of light. In Secs.~III and IV, we will investigate the deflection angles in the weak field approximation and in the strong field limit. In Sec.~V, we study the gravitational lens effects in the strong field limit in the Tangherlini spacetime. In Sec.~VI, we summarize our results. In this paper we use the units in which the light speed $c=1$ and Newton's constant $G=1$. ",
        "conclusions": "We discussed whether relativistic images by the strong field can be neglected in general. It is likely that a spacetime geometry with $1/r^n$ fall-off in the weak field is not unique but has many variants. The Tangherlini solution is one of the simplest models for our purpose. For recent investigations of exotic gravitational lenses, therefore, it would be of physical interest to examine Tangherlini lens both in the weak field and in the strong one.  We calculated the divergent part of the deflection angle for arbitrary $n$ and the regular part for $n=1$, $2$ and $4$ in the strong field limit, the deflection angle for arbitrary $n$ under the weak gravitational approximation and the relation between the size of the Einstein ring and the ones of the relativistic Einstein rings for arbitrary $n$. We showed that the relativistic images are always fainter than the images in the weak gravitational field.   We conclude that the images in the strong gravitational field have little effect on the total light curve and that the time-symmetric demagnification~\\cite{Abe_2010,Kitamura_Nakajima_Asada_2013} of the light curve will appear even after taking account of the images in the strong gravitational field for $n>1$."
    },
    "1402/1402.4017_arXiv.txt": {
        "abstract": "{It is very likely that most (perhaps all) Galactic globular clusters (GGCs) have experienced two or even more star-formation episodes. This is indicated, in particular, by peculiar chemical composition of the cluster stars which show large variation in the abundances of light elements, such as Li, C, N, O, Na, Mg, and Al. We studied the abundances of Na, Mg, and Al in the atmospheres of 103 red giant branch stars located below (49 stars) and above (54 stars) the RGB bump in the GGC 47 Tuc. Our results show that the spread of [Na/Fe] abundance ratios is about three times larger than that of [Mg/Fe]. Our data also confirm the existence of weak Na-Al correlation, similar to the one observed in other GGCs. At the same time, we find no evidence for the existence of three populations of stars characterized with different abundances of aluminum, as reported recently by \\citet{Ca13}. ",
        "introduction": "Galactic globular clusters (GGCs) are old, metal-poor stellar systems which for many years were thought to be good examples of simple stellar populations. However, recent photometric and spectroscopic observations seem to support the idea that stars in these stellar systems could have formed during two or more star formation episodes. This scenario is corroborated by the fact that abundances of light elements (Li, O, Na, Mg, Al) in the GGCs show various correlations/anti--correlations, such as Na--O anti--correlation \\citep{Kr94}, Li--O correlation \\citep{Pa05}, Li--Na anti--correlation \\citep{Bo07}, Mg--Al anti--correlation \\citep{Ca09}. It is important to mention that such abundance trends are not seen in the halo field stars of the same metallicity.  It is generally accepted today that stars characterized by different chemical composition belong to at least two different stellar generations: second generation stars (e.g., O-poor, Na-rich) are considered to be younger than first generation stars (O-rich, Na-poor), and therefore, must have formed from the matter enriched in chemical elements which were synthesized by the first generation stars \\citep[see, e.g.,][]{Gr01}. It is thought that two types of first generation stars could have been responsible for the enrichment of second generation stars: (i) AGB stars \\citep{Ve01}, or (ii) fast rotating massive stars \\citep{De07}. However, none of the two scenarios can fully account for the abundance trends observed in the GGCs. Therefore, further insights about the chemical evolution of the GGCs based on the results of homogeneous chemical analysis of large number of stars in multiple GGCs, as well as new theoretical interpretations, would be highly desirable. Galactic globular cluster 47 Tuc is particularly interesting target in this respect, due to its proximity and large angular diameter.  In this contribution we present first results of our study aimed to investigate the abundance trends of Na, Mg, and Al in 47 Tuc, to better constrain the possible chemical evolutionary scenarios of this GGC.   \\begin{figure} \\begin{center} \\includegraphics[scale=0.60]{Fig1.eps} \\caption{\\footnotesize $T\\rm_{eff}$--log \\textit{g} diagram of 47 Tuc showing RGB stars studied in this work.} \\label{HR} \\end{center} \\vspace{-5mm} \\end{figure} ",
        "conclusions": "The following mean abundance--to--iron ratios were obtained for 103 RGB stars located both below and above the RGB bump: $\\langle{\\rm [Na/Fe]}\\rangle=0.29\\pm0.13$ and $\\langle{\\rm [Mg/Fe]}\\rangle=0.32\\pm0.12$ (square root of the abundance dispersion in all individual stars is provided after the $\\pm$ sign). In case of Al, only stars below the RGB bump were studied so far (49 objects), resulting in $\\langle$[Al/Fe]$\\rangle$=0.28$\\pm$0.13.  We find that [Mg/Fe] abundance ratios do not show large star--to--star variation in the RGB stars of 47 Tuc but the spread in [Na/Fe] ratios is more than two times larger than that in [Mg/Fe] (Fig.~\\ref{abn}, \\textit{left panel}). Moreover, we found no clear anti-correlation neither between  Na and Mg, nor between Mg and Al (Fig.~\\ref{abn}). Nevertheless, it seems that our data suggests the existence of weak Na--Al correlation (Fig.~\\ref{abn}, \\textit{right}), similar to the one observed in other GGCs \\citep[see, e.g.,][]{Ca09}: the probability \\textit{p} that this correlation does not exist (estimated using Kendall's $\\tau$ test) is equal to $2.2 \\times {10^{-4}}$, correlation coefficient is $0.51$.  It may be interesting to note that stars located above the RGB bump may display the signs of an additional 'thermohaline mixing' event \\citep{Mu12}. However, we found no obvious differences in the abundance patterns of Na and Mg in our sample stars located below and above the RGB bump.  While the obtained results do not allow to unambiguously discriminate between the AGB and fast rotating massive star pollution scenarios, they do not seem to support the evidence of three distinct populations recently reported by \\citet{Ca13}."
    },
    "1402/1402.0919_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small The enormous radiative and mechanical luminosities of massive stars impact a vast range of scales and processes, from the reionization of the universe, to the evolution of galaxies, to the regulation of the interstellar medium, to the formation of star clusters, and even to the formation of planets around stars in such clusters. Two main classes of massive star formation theory are under active study, {\\it Core Accretion} and {\\it Competitive Accretion}. In Core Accretion, the initial conditions are self-gravitating, centrally concentrated cores that condense with a range of masses from the surrounding, fragmenting clump environment. They then undergo relatively ordered collapse via a central disk to form a single star or a small-$N$ multiple. In this case, the pre-stellar core mass function has a similar form to the stellar initial mass function. In Competitive Accretion, the material that forms a massive star is drawn more chaotically from a wider region of the clump without passing through a phase of being in a massive, coherent core. In this case, massive star formation must proceed hand in hand with star cluster formation. If stellar densities become very high near the cluster center, then collisions between stars may also help to form the most massive stars. We review recent theoretical and observational progress towards understanding massive star formation, considering physical and chemical processes, comparisons with low and intermediate-mass stars, and connections to star cluster formation.\\\\~\\\\~\\\\} ",
        "introduction": "}  Across the universe, massive stars play dominant roles in terms of their feedback and their synthesis and dispersal of heavy elements. Achieving a full theoretical understanding of massive star formation is thus an important goal of contemporary astrophysics. This effort can also be viewed as a major component of the development of a general theory of star formation that seeks to explain the birth of stars of all masses and from all varieties of star-forming environments.  Two main classes of theory are under active study, {\\it Core Accretion} and {\\it Competitive Accretion}. In Core Accretion, extending ``standard'' low-mass star formation theory ({\\em Shu et al.}, 1987), the initial conditions are self-gravitating, centrally concentrated {\\it cores} of gas that condense with a range of masses from a fragmenting {\\it clump} (i.e., protocluster) environment. These cores then undergo gravitational collapse via a central disk, to form a single star or small-$N$ multiple. The pre-stellar core (PSC) mass function (CMF) has a shape similar to the stellar initial mass function (IMF). In Competitive Accretion, gas that forms a massive star is drawn chaotically from a wider region of the clump, without ever being in a massive, coherent, gravitationally bound, starless core. Also, a forming massive star is always surrounded by a swarm of low-mass protostars. Competitive Accretion is sometimes said to lead naturally to the IMF ({\\it Bonnell et al.,} 2001; 2007): then the total mass of massive stars must be a small fraction of the total stellar mass formed from the clump. If the density of protostars congregating near the cluster center becomes sufficiently high, then stellar collisions may also assist in forming the most massive stars.  Recent advances in theoretical/numerical modeling of massive star formation involve inclusion of more physical processes, like radiation pressure, magnetic fields and protostellar outflows.  Observationally, progress has resulted from telescopes such as {\\it Spitzer}, {\\it Herschel}, {\\it SOFIA}, {\\it ALMA} and the {\\it VLA}.  Galactic plane surveys have yielded large samples of candidate massive protostars and their birth clouds.  This review aims to summarize massive star formation research, focusing on developments since the reviews of {\\em Beuther et al.} (2007), {\\em Zinnecker and Yorke} (2007) and {\\em McKee \\& Ostriker} (2007). We do not discuss formation of the first stars, which are thought to have been massive (e.g., {\\it Bromm}, 2013). Given the complexity of massive star formation, detailed comparison of theoretical predictions with observational results is needed for progress in understanding which accretion mechanism(s) is relevant and which physical and chemical processes are important. We thus first overview basic observed properties of massive star-forming regions (\\S\\ref{S:environ}), which set boundary conditions on theoretical models. Next we present a theoretical overview of physical processes likely involved in forming massive stars (\\S\\ref{S:theory}), including the different accretion models, protostellar evolution and feedback, and results from numerical simulations. We then focus on observational results on the earlier, i.e., initial condition (\\S\\ref{S:initial}) and later, i.e., accretion (\\S\\ref{S:accretion}) stages of massive star formation. Here we discuss astrochemical modeling, as well as general comparisons of massive star formation with intermediate/low-mass star formation.  The relation of massive star formation to star cluster formation is examined in \\S\\ref{S:cluster}. We conclude in \\S\\ref{S:conclusions}.   \\subsection{The Birth Environments of Massive Stars}\\label{S:environ}  The basic physical properties of regions observed to be forming or have formed massive stars, i.e., gas clumps and young star clusters, are shown in Fig.~\\ref{fig:overview}, plotting mass surface density, $\\Sigma = M/(\\pi R^2)$, of the structure versus its mass, $M$.  Stars, including massive stars, form in molecular gas, that is mostly found in giant molecular clouds (GMCs) with $\\Sigma\\sim 0.02\\:\\gcc$. Note $1.0\\:\\gcc \\equiv 4790\\:M_\\odot\\:{\\rm pc^{-2}}$, for which $N_{\\rm H} = 4.27\\times 10^{23}\\:{\\rm cm^{-2}}$ (assuming $n_{\\rm He}=0.1 n_{\\rm H}$ so mass per H is $\\mu_{\\rm H}=2.34\\times 10^{-24}\\:{\\rm g}$) and visual extinction is $A_V=(N_{\\rm H}/2.0\\times 10^{21}$~cm$^{-2})$~mag $= 214$~mag. However, star formation is seen to be localized within star-forming clumps within GMCs, which typically have $\\Sigma_{\\rm cl}\\sim 0.1 - 1\\:\\gcc$. Some massive systems, usually already-formed star clusters, have $\\Sigma$ up to $\\sim 30\\:\\gcc$.   In terms of $\\Sigma_{\\rm cl}$ (in $\\rm g\\; cm^{-2}$) and $M_{\\rm cl,3}=M_{\\rm cl}/(1000\\:M_\\odot)$, the radius and (H number) density of a spherical clump are \\beqa R_{\\rm cl}&=&0.258 M_{\\rm cl,3}^{1/2} \\Sigma_{\\rm cl}^{-1/2}~~\\pc,\\label{eq:rcl}\\\\ \\bar{n}_{\\rm H,cl}&=&4.03\\times 10^5 \\Sigma_{\\rm cl}^{3/2} M_{\\rm cl,3}^{-1/2}~~~\\mbox{cm\\eee}. \\eeqa Gas clumps massive enough to form a cluster of mass $M_{\\rm *cl}\\sim 500\\:M_\\odot$, i.e., with median expected maximum stellar mass $\\sim30\\:M_\\odot$ (for Salpeter IMF from 0.1 to 120~$M_\\odot$), are thus $\\sim$0.3~pc in size (if $\\Sigma_{\\rm cl}\\sim 1\\:{\\rm g\\:cm^{-2}}$ and efficiency $\\epsilon_{\\rm *cl}\\equiv M_{\\rm *cl}/M_{\\rm cl}\\sim 0.5$), only moderately larger than the $\\sim$0.1~pc sizes of well-studied low-mass starless cores in regions such as Taurus ({\\it Bergin and Tafalla,} 2007). However, mean densities in such clumps are at least ten times larger.    \\begin{figure*} \\epsscale{2.05} \\plotone{figppvi.pdf} \\caption{\\small \\label{fig:overview} The Environments of Massive Star Formation.  Mass surface density, $\\Sigma\\equiv M/(\\pi R^2)$, is plotted versus mass, $M$. Dotted lines of constant radius, $R$, H number density, $n_{\\rm H}$ (or free-fall time, $t_{\\rm ff}= (3\\pi/[32G\\rho])^{1/2}$), and escape speed, $v_{\\rm esc} = (10/\\alpha_{\\rm vir})^{1/2}\\sigma$, are shown. Stars form from molecular gas, which in the Galaxy is mostly organized into GMCs. Typical $\\rm ^{12}CO$-defined GMCs have $\\Sigma\\sim 100\\:M_\\odot\\:{\\rm pc^{-2}}$ ({\\it Solomon et al.,} 1987) (see {\\it Tan et al.,} 2013a for detailed discussion of the methods for estimating $\\Sigma$ for the objects plotted here), although denser examples have been found in Henize 2-10 ({\\it Santangelo et al.,} 2009). The $\\rm ^{13}CO$-defined clouds of {\\it Roman-Duval et al.} (2010) are indicated, along with $\\rm HCO^+$ clumps of {\\it Barnes et al.,} (2011), including G286.21+0.17 ({\\it Barnes et al.,} 2010). Along with G286, the BGPS clumps ({\\it Ginsburg et al.,} 2012) and the Galactic Center ``Brick'' ({\\it Longmore et al.,} 2012) are some of the most massive, high-$\\Sigma$ gas clumps known in the Milky Way. Ten example Infrared Dark Clouds (IRDCs) ({\\it Kainulainen and Tan,} 2013) and their internal core/clumps ({\\it Butler and Tan,} 2012) are shown, including the massive, monolithic, highly-deuterated core C1-S ({\\it Tan et al.,} 2013b). CygX-N63, a core with similar mass and size as C1-S, appears to be forming a single massive protostar ({\\it Bontemps et al.,} 2010; {\\it Duarte-Cabral et al.,} 2013). The IRDC core/clumps overlap with Massive Star-Forming (MSF) core/clumps ({\\it Mueller et al.,} 2002).  Clumps may give rise to young star clusters, like the ONC (e.g., {\\it Da Rio et al.,} 2012) and NGC 3603 ({\\it Pang et al.,} 2013) (radial structure is shown from core to half-mass, $R_{1/2}$, to outer radius), or even more massive examples, e.g., Westerlund 1 ({\\it Lim et al.,} 2013), Arches ({\\it Habibi et al.,} 2013), Quintuplet ({\\it Hu{\\ss}mann et al.,}~2012) (shown at $R_{1/2}$), that are in the regime of ``super star clusters'' (SSCs), i.e., with $M_*\\gtrsim 10^4\\:M_\\odot$. Example SSCs in the Large Magellanic Cloud (LMC) (R136, {\\it Andersen et al.,} 2009) and Small Magellanic Cloud (SMC) (NGC 346, {\\it Sabbi et al.,} 2008) display a wide range of $\\Sigma$, but no evidence of IMF variation (\\S\\ref{S:IMF}).  Even more massive clusters can be found in some dwarf irregular galaxies, such as NGC 1569 ({\\it Larsen et al.,} 2008) and NGC 5253 ({\\it Turner and Beck,} 2004), and starburst galaxy M82 ({\\it McCrady and Graham,} 2007).  } \\end{figure*} ",
        "conclusions": "}\\label{S:conclusions}  It is a challenge to understand the wide variety of interlocking physical and chemical processes involved in massive star formation. Still, significant theoretical progress is being made in modeling these processes, both individually and in combination in numerical simulations. However, such simulations still face great challenges in being able to adequately resolve the scales and processes that may be important, including MHD-driven outflows, radiative feedback and astrochemistry. There are also uncertainties in how to initialize these simulations. Accurate prediction of the IMF, including massive stars, of a cluster forming under given environmental conditions remains a distant goal.  Close interaction with observational constraints is essential. Here rapid progress is also being made and, with the advent of {\\it ALMA}, this should only accelerate. One challenge is development of the astrochemical sophistication needed to decipher the rich variety of diagnostic tracers becoming available for both pre- and protostellar phases. Determination of pre-stellar core mass functions and resolution of massive protostellar accretion disks, including possible binary formation, and outflows are important goals.  Core and Competitive Accretion theories are being tested by both simulation and observation. Core Accretion faces challenges of understanding fragmentation properties of magnetized, turbulent gas, following development of accretion disks and outflows from collapsing cores, and assessing the importance of external interactions in crowded cluster environments. Competitive Accretion is also challenged by theoretical implementation of realistic feedback from MHD outflows and by observations of massive starless cores, together with apparent continuities and similarities of the star formation process across protostellar mass and luminosity distributions. There is much work to be done!    \\textbf{Acknowledgments.} We thank S. Kong, A. Myers, \\'A. S\\'anchez-Monge, Y. Zhang for figures and discussions, and H. Beuther, M. Butler, P. Clark, N. Da Rio, M. Hoare, P. Klaassen, P. Kroupa, W. Lim \\& a referee for discussions."
    },
    "1402/1402.0128_arXiv.txt": {
        "abstract": "With six recorded nova outbursts, the prototypical recurrent nova T Pyxidis is the ideal cataclysmic variable system to assess the net change of the white dwarf mass within a nova cycle. Recent estimates of the mass ejected in the 2011 outburst ranged from a few $\\sim 10^{-5}M_{\\odot}$ to $3.3 \\times 10^{-4}M_{\\odot}$, and assuming a mass accretion rate of $10^{-8}-10^{-7}M_{\\odot}$yr$^{-1}$ for 44yrs, it has been concluded that the white dwaf  in T Pyx is actually losing mass. Using NLTE disk modeling spectra to fit our recently obtained Hubble Space Telescope (HST) COS \\& STIS spectra, we find a mass accretion rate of up to two orders of magnitude larger than previously estimated. Our larger mass accretion rate is due mainly to the newly derived distance of T Pyx (4.8kpc, larger than the previous 3.5kpc estimate), our derived reddening of E(B-V)=0.35 (based on combined IUE and GALEX spectra) and NLTE disk modeling (compared to black body and raw flux estimates in earlier works). We find that for most values of the reddening ($0.25 \\le E(B-V) \\le 0.50$) and white dwaf  mass ($0.70 M_{\\odot} \\le M_{wd} \\le 1.35M_{\\odot}$) the accreted mass is larger than the ejected mass. Only for a low reddening ($\\sim 0.25$ and smaller)  {\\it combined with} a large white dwaf mass ($0.9 M_{\\odot}$ and larger) is the ejected mass larger than the accreted one. However, the best results are obtained for a larger value of the reddening. ",
        "introduction": "Cataclysmic variables (CVs) are binaries in which a white dwarf (WD; the primary) accretes hydrogen-rich material from a Roche-lobe filling secondary star. When sufficient matter is accreted, it undergoes a thermonuclear runaway (TNR):  the classical nova eruption. The TNR ignition starts at the base of the accreted shell when the pressure and temperature are high enough for the CNO cycle burning of hydrogen \\citep{pac65,sta72}. CVs with recorded multiple nova outbursts are classified as recurrent novae (RNe). In RNe, a large mass WD and a high mass accretion rate are believed to lead to the short recurrence time between nova outbursts.  T Pyxidis, the prototypical RN, is known to have erupted \\citep{web87,sch13} in 1890, 1902, 1920, 1944 and 1967 - roughly every $\\sim$20 yr. It erupted again in April 2011 \\citep{waa11}, after a lapse of 44 years. The mass accreted between eruptions and the mass ejected during eruptions determine the net change in the WD mass. If the WD mass increases, then eventually it could  reach the Chandrasekhar limit and explode as a type Ia supernova (SN Ia; \\citet{whe73,ibe84}). Due to its relatively frequent       outbursts, T Pyx offers a promising possiblity of assessing the WD mass change within a nova cycle. Recent works \\citep{sel08,pat13,nel12} have claimed that the WD in T Pyx is {\\it losing} mass and will never become a SN Ia. In this Letter, we present a synthetic spectroscopic analysis of the first far ultraviolet (FUV; 900\\AA - 1700\\AA ) spectra of T Pyx obtained with Hubble. Our objective is to estimate the rate of mass accretion following its July 2011 outburst and compare it with estimates of the mass accretion rate in the quiescent interval preceding the July 2011 outburst. ",
        "conclusions": "Our results are recapitulated in Fig.7, where we draw (solid lines) the mass accreted over 44 years as computed from our model fittings, for 3 different values of the reddening.  We also draw (dotted line) the lower limit for the mass accreted derived from integrating the UV flux over the wavelengths 900\\AA\\ to 3200\\AA\\ $$ \\Delta M_{min} = 44 \\times \\dot{M}_{min} = 44 \\times \\int F_{\\lambda} d\\lambda ~. $$ Since the disk at high mass accretion rate emits only a fraction of its energy in the UV (and a much smaller fraction is emitted in the optical), the integrated UV flux is clearly a lower limit for the mass accretion rate. We draw this lower limit for three different values of the reddening as well. For comparison we draw the ejecta mass as derived by radio observations by \\citet{nel12} (denoted with a square, \\#1), who assumed $M_{wd}=0.7M_{\\odot}$ and E(B-V)=0.50. This ejecta mass is to be compared with our accreted mass derived from our disk modeling assuming E(B-V)=0.50. Our results show that the accreted envelope (for the particular values of the parameters) is about 4 times larger than the mass ejected as estimated by \\citet{nel12}. On the graph, we also show the mass of the ejected envelope as estimated by \\citet{pat13}  assuming that energy and momentum are conserved (dashed line, \\#2).  We see that the disk model fittings imply that more mass is accreted than ejected, except when assuming a reddening of E(B-V)=0.25 {\\it in combination with} a WD          mass $M_{wd}> 0.9M_{\\odot}$.   However, for this value of the reddening, the disk model (for all WD masses considered) did not fit the observed spectrum, indicating that the reddening might indeed be larger (i.e. 0.35).  We note that the $\\sim 10$\\% error in the distance $4.8 \\pm 0.5 $kpc translates into an error of the order of 20\\% in the mass accretion rate, which does not change our results significantly. We also emphasize that the difference between our results and previously derived mass accretion rates is due mainly to the now larger distance estimates (4.8kpc vs 3.5kpc) as well as a larger reddening (0.35 against 0.25).  Our      results are: (i) the determination of the long-time-baseline accretion rate in T Pyx, using for the first time realistic disk models, (ii) the demonstration that the accretion rate in T Pyx has returned to exactly its pre-outburst UV spectral energy distribution. This accretion rate, $\\sim 10^{-6}M_{\\odot}$yr$^{-1}$ (and higher for increasing values of the reddening), is much higher than previously estimated.  It is possible that T Pyx sustains such a high mass accretion rate because of irradiation of the secondary star. The irradiation of the donor star by the extremely hot WD and inner disk will not only ``puff up'' the donor's radius but may also drive a wind off the secondary \\citep{kni00}. Since the orbital period and binary separation of T Pyx are smaller than all other known recurrent novae, the effects of irradiation may be more pronounced.  Several different authors reported that quasi-static evolutionary sequences of accreting WDs expand in to red giant structures at accretion rates exceeding $10^{-7} M_{\\odot}$yr$^{-1}$ \\citep{pac78,sio79,ibe82} but other evolutionary accreting WD          sequences followed hydrodynamically through multiple nova outbursts (e.g. \\citet{pri95,yar05,sta12,ida13}) reveal thermonuclear shell flashes at these high rates."
    },
    "1402/1402.5227_arXiv.txt": {
        "abstract": "We present a sample of 33 spectroscopically confirmed z $\\sim$ 3.1 \\lya\\ emitting galaxies (LAEs) in the COSMOS field.  This paper details the narrowband survey we conducted to detect the LAE sample, the optical spectroscopy we performed to confirm the nature of these LAEs, and a new near-infrared spectroscopic detection of the \\oiii\\ 5007 \\AA\\ line in one of these LAEs.  This detection is in addition to two \\oiii\\ detections in two z $\\sim$ 3.1 LAEs we have reported on previously \\citet{mcl11}.    The bulk of the paper  then presents detailed constraints on the physical characteristics of the entire LAE sample from spectral energy distribution (SED) fitting.  These characteristics include mass, age, star-formation history, dust content and metallicity.  We also detail an approach to account for nebular emission lines in the SED fitting process - wherein our models predict the strength of the \\oiii\\ line in an LAE spectrum.  We are able to study the success of this prediction because we can compare the model predictions to our actual near-infrared observations both in galaxies that have \\oiii\\ detections and those that yielded non-detections.  We  find a median stellar mass of 6.9 $\\times$ 10$^{8}$ \\msol\\ and a median star formation rate weighted stellar population age of 4.5 $\\times$ 10$^{6}$ years.  In addition to SED fitting, we quantify the velocity offset between the \\oiii\\ and \\lya\\ lines in the galaxy with the new \\oiii\\ detection, finding that the \\lya\\ line is shifted 52 \\kms\\ red-ward of the \\oiii\\ line, which defines the systemic velocity of the galaxy. ",
        "introduction": "High redshift Lyman alpha emitting galaxies (LAEs) are now routinely identified via narrowband detection methods (e.g. Cowie \\& Hu, 1998, Malhotra \\& Rhoads, 2002, 2004, Ouchi et al. 2003, Gawiser et al. 2006).  Now that samples of these galaxies can be more easily compiled at a variety of redshifts, attention has turned to deriving the physical characteristics of these galaxies by fitting spectral energy distributions (SEDs) to their observed photometry of these galaxies \\citep{gaw06, gaw07, pirz07, nil07, nil11, fink07, fink08, fink09, fink11b, lai07, lai08, ono10a, acq12}.  The majority of early work in SED fitting (e.g. Gawiser et al. 2006, Finkelstein et al. 2007, Nilsson et al. 2007)  relied on deriving average LAE characteristics from stacked LAE samples, but stacked analyses may not reveal the full distribution of LAE characteristics. Most efforts to date have found LAEs to be largely young or of intermediate ages and having characteristically small masses \\citep{pirz07, gaw07, fink09, cow11}, but SED fitting procedures tend to vary from author to author, making direct comparisons of derived characteristics difficult from sample to sample.  In addition, SED fitting procedures for starbursting galaxies have been evolving recently to account for contamination of observed photometry from  nebular emission lines. \\citet{sd09} and others have demonstrated that failure to include these lines, produced from hot gas in star forming regions, can drastically alter the ages and masses derived from SED fitting.  In this paper we present a simple way to account for nebular emission during SED fitting in our sample of 33 spectroscopically confirmed z $\\sim$ 3.1 LAEs.  The technique we outline in this paper allows us to predict the strength of the \\oiii\\ nebular emission line, which we can compare to the NIR detections and upper limits we have made of this line in six z $\\sim$ 3.1 LAEs.  In Section \\ref{sec:obsdat} we present the extensive observations that form the foundation of this paper, including a narrowband survey to find LAE candidates, optical spectroscopy to confirm LAE candidates and NIR spectroscopy to look for rest-frame optical nebular emission lines in these LAEs.  We also present our data reduction techniques in this section.  In Section \\ref{sec:res1} we present our results from optical and NIR spectroscopy, including a new \\oiii\\ measurement in one LAE and the subsequent velocity offset between \\oiii\\ and \\lya\\ that we measure in this object.  Section \\ref{sec:sed} outlines our methods for SED fitting, including the introduction of a new method to account for nebular emission lines in the SED fitting process.  We present our results from SED fitting in Section \\ref{sec:res2}.  Finally, in Section \\ref{sec:disc} we discuss the ability of our SED fitting process to match our observations of the \\oiii\\ line in LAEs.  We also compare our SED results to those presented by other authors.  Where relevant, we adopt the standard cosmological parameters H$_{0} =$ 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = $ 0.3, and $\\Omega_{\\Lambda} =$ 0.7 \\citep{sper}.  Also we use the following vacuum wavelengths, 1215.67  \\AA\\ for \\lya,  3727.092/3729.875 \\AA\\ for [O\\,{\\sc ii}], 4862.683 \\AA\\ for H$\\beta$  and 4960.295/5008.240 for [O\\,{\\sc iii}] from the Atomic Line List v2.04\\footnote{http://www.pa.uky.edu/$\\sim$peter/atomic/index.html}.  All quoted magnitudes are AB magnitudes. ",
        "conclusions": "We have presented one new \\oiii\\ detection in a z $\\sim$ 3.1 LAE.  Combining this new detection with the two we presented in Mc11, we are able to present a total of three measurements of the velocity offset between \\lya\\ and \\oiii\\ in these z $\\sim$ 3.1 LAEs, ranging from 52 -- 342 \\kms.  This new result is still consistent with the outflow models explored in Mc11.  In addition to the new \\oiii\\ detection, we have put forth a simple method to account for nebular emission in high-z starbursting galaxies, motivated by our three \\oiii\\ measurements.  We have individually fit 33 z $\\sim$ 3.1 LAEs using this powerful yet simple method to account for nebular emission line contributions to galaxy SEDs.  From these fits we find constraints on age, mass, dust content, metallicity, star formation history, and \\oiii\\ line flux.  We find that our sample has quite diverse characteristics, but some generalizations can be made.  For instance, a majority of the galaxies are fit with a single instantaneous burst or exponentially decreasing star formation history.  As a whole, the sample has only moderate amounts of dust, and sub-solar metallicity.  Mass and age solutions vary widely, but median values of 4.5 $\\times$ 10$^{6}$ years (full range is 1.4 $\\times$ 10$^{6}$ -- 4.6 $\\times$ 10$^{8}$ years)  and 6.9 $\\times$ 10$^{8}$ \\msol\\ (full range is 7.1 $\\times$ 10$^{7}$ -- 6 $\\times$ 10$^{10}$ \\msol)  are found.  Finally, most of the galaxies are best fit with an \\oiii\\ line contributing additional flux to the \\kfil\\ band, with an average flux of 7.3 $\\times$ 10$^{-17}$ \\lf\\ (or an average of 9.7 $\\times$ 10$^{-17}$ \\lf\\ among the 25 galaxies with non-zero line flux solutions, ranging from 7.0 $\\times$ 10$^{-18}$ -- 3.7 $\\times$ 10$^{-16}$ \\lf).  The \\oiii\\ line strength predictions from our new SED fitting methodology have reasonably matched the observations of the \\oiii\\ line in the six objects for which we can make this comparison.   These predictions gives us confidence that these results can be used to select the LAEs mostly likely to yield \\oiii\\ detections in future NIR observations and aids in planning adequate integration times for the most efficient use of such future NIR observing time.  Further observations of LAEs in the NIR will allow us to fill in the distribution of velocity offsets found in LAEs at this redshift, and will allow us to further test the validity of \\oiii\\ line strength predictions from our SED fitting process.  In the meantime we have, with this work, provided a comprehensive picture of LAE characteristics in a large sample of individually examined objects."
    },
    "1402/1402.3020.txt": {
        "abstract": "We propose a catastrophic eruption model for magnetar's enormous energy release during giant flares, in which a toroidal and helically twisted flux rope is embedded within a force-free magnetosphere. The flux rope stays in stable equilibrium states initially and evolves quasi-statically. Upon the loss of equilibrium point is reached, the flux rope cannot sustain the stable equilibrium states and erupts catastrophically. During the process, the magnetic energy stored in the magnetosphere is rapidly released as the result of destabilization of global magnetic topology. The magnetospheric energy that could be accumulated is of vital importance for the outbursts of magnetars. We carefully establish the fully open fields and partially open fields for various boundary conditions at the magnetar surface and study the relevant energy thresholds. By investigating the magnetic energy accumulated at the critical catastrophic point, we find that it is possible to drive fully open eruptions for dipole dominated background fields. Nevertheless, it is hard to generate fully open magnetic eruptions for multipolar background fields. Given the observational importance of the multipolar magnetic fields in the vicinity of the magnetar surface, it would be worthwhile to explore the possibility of the alternative eruption approach in multipolar background fields. Fortunately, we find that flux ropes may give rise to partially open eruptions in the multipolar fields, which involve only partial opening up of background fields. The energy release fractions are greater for cases with central-arcaded multipoles than those with central-caved multipoles emerged in background fields. Eruptions would fail only when the centrally-caved multipoles become extremely strong. ",
        "introduction": "Two small subsets of neutron stars $-$ anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are interpreted as magnetars, which are neutron stars endowed with ultra-strong ($10^{14}-10^{15}$G) magnetic fields \\citep{Maze79,MS95,Kouv98,GKW02}. The magnetic energy dissipation is commonly believed to account for their high energy persistent emissions and spasmodic outbursts \\citep{DT92,WT06,Mere08}. One of the most intriguing phenomena related to magnetars is the episodic giant flare, which involves tremendous magnetic energy release \\citep{Hurl05}. The physical origin of giant flares still remains the bewildering enigma in high energy astrophysics. The magnetospheric eruption models \\citep[e.g.,][]{Lyut06,GH10,Yu12,YH13} can naturally explain the short timescale, $\\sim0.25$ms, of the rise time in giant flares \\citep{Palm05}, although the precise mechanism of the eruption is still under debate \\citep{Mere13}. It is worthwhile to note that, in the magnetospheric eruption model, the magnetosphere stays in a stable equilibrium state at the pre-eruptive stage, in which the magnetosphere evolves quasi-statically. As a result, the magnetic energy released in an eruption is gradually accumulated on a timescale much longer the dynamical timescale of giant flares. Since the energy accumulation takes place very gradually, the question how such long timescale events could initiate the sudden energy release, i.e., giant flares, on a very short dynamical timescale, remains elusive for the magnetospheric eruption model.  To resolve this puzzle, a catastrophic flux rope eruption model has been put forward to explain the magnetar giant flare \\citep{Yu12,YH13}. In this particular scenario, the giant flare is considered to be driven by the destabilization of large-scale magnetospheric magnetic fields rather than the abrupt fracture of the neutron star crust. The most distinguishing feature in our model is that a helically twisted flux rope is embedded within the magnetosphere. Magnetic flux ropes could be naturally generated due to the the magnetic helicity injections from the magnetar interior (Thompson et al .2002; Lyutikov 2006; G\\\"{o}tz et al. 2007; Gill \\& Heyl 2010). Such flux ropes are also an indispensable ingredient to explain the radio afterglow of SGR1806 \\citep{Gaen05,Lyut06}. It is interesting to note that the interior of the flux rope is helically twisted. The magnetic twist is locally confined within the flux rope, which is at variance with the global twist proposed by recent authors \\citep{Lyut13,Parf13}. These locally twisted magnetic features, when compared to the global twist, seem to be more relevant to the recent observations \\citep[e.g.,][]{Wood07,PG08}.  The flux rope eruption model tries to resolve a primary issue concerning the trigger mechanism of the magnetar eruption. The most appealing characteristic is that, during the flux rope's evolution, it could make the catastrophic state transition from stable equilibrium states to unstable equilibrium states spontaneously in accordance with the variations at the neutron star surface. The emergence of flux from the interior \\citep{KR98,GMH07} and/or the shuffling of the crust \\citep{Rude91} causes the flux rope to evolve to a critical loss of equilibrium point, beyond which no stable equilibrium state can be sustained and the flux rope erupts abruptly.  It is widely accepted that the energy is progressively accumulated in an initially closed force-free field before the flux rope reaches the critical loss of equilibrium point. The flux rope's subsequent catastrophic eruption, beyond this particular point, leads to the opening up of initially closed magnetic field configuration as well as huge energy release. Since the stored energy prior to an eruption is determined by the intrinsic properties of the magnetosphere rather than the tensile strength of the crust \\citep{Yu12}, a fundamental question is raised naturally as to whether the magnetic energy in the force free magnetosphere could build up enough energy to support eruptive giant flares.  The fierce eruptive event is thought to open up the pre-eruptive, originally closed magnetic field lines. Observationally, in the post-eruption epoch of giant flares, magnetic field configurations are indeed inferred to stretch outwards to form an open field structure \\citep[e.g.][]{Wood01}. From the perspective of energy budget, a closed field configuration is physically favored for giant flares only if the magnetic energy stored in magnetosphere can exceeds that stored in an open field configuration. In other words, in addition to power the giant flares, the magnetic free energy has to be able to open up the initially closed magnetic fields. This constitutes a serious bottleneck for magnetar giant flares because it is realized that to open up the initially closed field lines requires considerable extra work to be done \\citep{Aly84,Aly91,Stur91,Yu12}. To get over the energy threshold constrained by the post-eruptive open field configurations, the pre-eruptive magnetosphere must accumulate more energy in excess of the threshold set by the open field configurations. It should be pointed out that the field lines can also be opened up at a greater distance from the star by strong neutron star winds \\citep[e.g.,][]{Bucc06}. How the wind affect the magnetar eruption energetics is still an open issue. For simplicity, we only consider the case in which the field lines are opened up by the process of magnetic flux injection from the neutron star interior.  The observational feature of strong four-peaked pattern in the pulse profile of the 1998 August 27 event from SGR 1900+14 indicates that the geometry of the magnetic field was quite complicated in regions close to the star \\citep{Fero01}. Recent calculations also show that multipolar magnetic fields may also have important effects on the emission of the magnetars \\citep{Pava09}. As a result, it is reasonable to infer that, in the very vicinity of the magnetar surface, the field configuration involves higher multipoles. The electric currents formed during the birth of magnetars slowly push out from within the magnetar and generate active regions on the magnetar surface. These active regions manifest themselves as the multipolar regions on the magnetar surface. The flux rope eruption in multipolar background fields, unlike the behavior in dipole background fields, may just involve opening-up of part of the closed magnetic flux systems. Thus, it is interesting to explore the physical behavior of the flux rope in response to these more complex boundary conditions. However, for more complex boundary conditions, no solid investigations about the energy threshold specified by the fully/partially open field configurations have been performed. A related question, which is of crucial significance for the physical feasibility of the flux rope eruption model, i.e., whether the flux rope could build up enough energy to drive the giant flare with these complex boundary conditions, also remains to be answered.  In this paper, we will establish a force-free magnetosphere model with a helically twisted flux rope and examine the physical response of the flux rope to the variations of the background multi-polar magnetic fields. We perform rigorous calculations on the magnetic energy accumulation in magnetospheres in various configurations. Specifically, boundary conditions containing the dipolar term and the high order multipolar term are considered in the background closed magnetic field. Two kinds of open configurations, the partially open and fully open fields, are considered to provide energy thresholds for eruptions. This paper is structured as follows. The model of pre-eruptive magnetospheres with flux ropes and the multipolar boundary conditions are introduced in Section 2. The physical behavior of the pre-eruptive flux rope is described in Section 3, including the equilibrium constraints and catastrophic loss of equilibrium of the flux rope. The energetics of the flux rope in the multipolar background fields are investigated in Section 4. Conclusions and discussions are provided in Section 5.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We propose a force-free magnetospheric model with an embedded helically twisted flux rope. With the gradual variations at the magnetar surface, the flux rope evolves quasi-statically in stable equilibrium states. Upon the loss of equilibrium point is reached, the global magnetospher is then destabilized and the flux rope erupts catastrophically. During the process, the original closed flux systems would be opened up, accompanied by rapid release of the magnetic energy stored in magnetosphere. This energy release is of vital importance for the outbursts of magnetars. However, the feasibility that the flux systems' opening-up could be achieved depends on whether the amount of energy accumulated prior to the flux rope eruptions could surpass the energy thresholds constrained by the post-eruptive open magnetic topologies.  In this paper, we adopt boundary conditions, which include both the contribution from a dipolar component and a high order multipolar component, to illustrate the complicated geometry of magnetic field close to magnetar surface. We establish fully open field for the dipole dominated magnetic fields, which involves the opening up the single closed flux system in the backgrounds. For multipole dominated closed background, we establish the partially open field, which involves opening up of part of the closed flux systems, as well as the fully open field, which involves opening up of all closed flux systems. The opening up of closed magnetic fields requires certain amount of work to be done to overcome the attractive magnetic tension force. Since partially open fields require less closed flux systems to be opened up, the energy thresholds constrained by the partially open fields is lower than those by the fully open fields. Both the field configuration and the magnetic energy threshold of the two kinds of open field are examined. Then we carefully investigate the magnetic energy accumulation process before the catastrophe, especially the magnetic energy stored at the critical catastrophic point.  We find that it is possible to fully open up dipole dominated background fields for catastrophic eruptions of flux ropes. However, it is generally difficult to fully open up multipole dominated background fields. In most cases with multipole dominated backgrounds, the magnetic energy stored at critical pre-eruptive point is significantly lower than the fully open thresholds, which suggests that the flux rope can not support fully open eruptions. Fortunately, we find that the accumulated magnetic energy at the critical point is higher than the partially open thresholds. This provides an alternative opportunity for the flux rope to erupt in the multipolar magnetosphere. The multipole dominated fields can be either centrally-caved or centrally-arcaded, depending on the flux profiles on the magnetar surface. Generally speaking, the magnetic energy stored in critical pre-eruptive magnetosphere surpasses the partially open energy threshold about $10\\%\\sim25\\%$, if the flux rope is initially embedded in a centrally-arcaded background field. For a flux rope initially embedded in a centrally-caved background field, the magnetic energy stored in critical pre-eruptive magnetosphere could surpass the partially open threshold, if the multipolar component is mildly strong, i.e., $1/6<a_1\\lesssim0.75$. The energy release fraction is within $\\sim5\\%$. If the multipolar component becomes even stronger, $a_1\\gtrsim0.75$, the accumulated magnetic energy cannot go beyond the partially open threshold and the partially open eruption of flux rope in not possible.  The magnetic energy of the critical pre-eruptive state in excess of the fully or partially open threshold is assumed to be released in fast dynamical timescale. Observationally, the total magnetic energy in the magnetosphere is about $\\sim 10^{46} ({\\rm B}/10^{14}G)^2 (r_s/10{\\rm km})^3$ergs, the giant flare is typically $10^{44}$ergs, so only $1\\%$ of magnetic energy release in the magnetosphere could account for the giant flares. Theoretically, all cases with surplus energy fraction larger than $1\\%$ are possible to drive magnetar giant flares. Specifically for boundary conditions adopted in this paper, most cases with $-1\\le a_1\\lesssim0.75$ are possible to drive magnetar giant flares.  In addition to the physical processes considered in this paper, the magnetic field can also be opened up by the strong neutron star wind (Bucciantini et al. 2006). If the field lines are opened up by the neutron star wind at a larger distance from the neutron star, the amount of open magnetic flux would be reduced, and the magnetic energy to support the eruption would also decrease. To study how the neutron star wind affects the flux rope eruption energetics, we need to establish a model with a current sheet in the magnetosphere. Currently, we are now trying to construct a magnetosphere model to account for this important process. We leave the investigation about the situations with current sheet formation in a companion paper \\citep{HYprep}.    \\begin{figure} \\includegraphics[scale=0.9]{f1.eps} \\caption{\\label{M0} {\\it Left}: Equilibrium curve of the flux rope in a dipole-dominated background with $a_1=-1/4$ and $r_{00}=0.01$, which shows the dependence of the height of the flux rope on the magnetic flux at the magnetar surface. The magnetospheric field configuration . The lower stable branch and the upper unstable branch are shown in solid line and dotted line, respectively. The red dot marks the critical point. {\\it Right}: Magnetosphere with an embedded flux rope in the pre-eruptive critical state. The thick solid semi-circle represents the magnetar surface. The dashed line represents a circle with a radius of the critical height of the flux rope $h_c$. } \\end{figure}  \\begin{figure} \\includegraphics[scale=0.9]{f2.eps} \\caption{\\label{M1d5} The same as Fig.\\ref{M0} but with a multipole-dominated  background with $a_1=2/3$. } \\end{figure}  \\begin{figure} \\includegraphics[scale=0.9]{f3.eps} \\caption{\\label{W0} {\\it Left}: The energy accumulation process before the catastrophe, $W_\\mathrm{pre}/W_\\mathrm{pot}$, in a background field with $a_1=-1/4$, along the stable branch as shown in Fig. \\ref{M0}. The fully open threshold is shown in dashed-three-dotted line. The accumulated magnetic energy goes beyond the fully open field threshold. {\\it Right}: The energy accumulation before the catastrophe in a background field with $a_1=2/3$, along the stable branch as shown in Fig. \\ref{M1d5}. The fully open and partially open thresholds are shown in dashed-three-dotted and dashed-dotted lines, respectively. The accumulated magnetic energy goes beyond the partially open field threshold, but still lower than the the fully open field threshold. } \\end{figure}  \\begin{figure} \\includegraphics[scale=0.9]{f4a.eps} \\includegraphics[scale=0.9]{f4b.eps} \\caption{\\label{WWpot1} {\\it Left}: Accumulated pre-eruptive magnetic energy as functions of $r_{00}$ for centrally-caved field with $a_1=0.2$ (black), $a_1=1/3$ (blue), and $a_1=1$ (red). The corresponding threshold values of partially open fields are shown in dashed-dotted lines in the same color. The corresponding threshold values of fully open fields are shown in dashed-three-dotted lines in the same color. {\\it Right}: The same as {\\it left} panel but for centrally-arcaded field with $a_1=-0.2$ (black), $-1/3$ (blue), and $-1$ (red). } \\end{figure}  \\begin{figure} \\includegraphics[scale=0.9]{f5.eps} \\caption{\\label{WWpot2} Accumulated pre-eruptive magnetic energy as functions of $a_1$ with $r_{00}=0.01$ (black), $0.05$ (grey), and $0.1$ (light grey). The corresponding threshold values of partially-open fields and fully-open fields are shown in dashed-dotted line and dashed-three-dotted line, respectively. Two vertical lines represent the separatrix between the dipole-dominated fields and the multipole-dominated fields. } \\end{figure}    %\\newpage  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1402/1402.0291_arXiv.txt": {
        "abstract": "The bright gamma-ray quasar 4C +55.17 is a distant source ($z = 0.896$) with a hard spectrum at GeV energies as observed by the Large Area Telescope (LAT) on board the {{\\it Fermi}} satellite. This source is identified as a good source candidate for very-high-energy (VHE; $> 30$ GeV) gamma rays. In general VHE gamma rays from distant sources provide an unique opportunity to study the extragalactic background light (EBL) and underlying astrophysics. The flux intensity of this source in the VHE range is investigated. Then, constraints on the EBL are derived from the attenuation of gamma-ray photons coming from the distant blazar. We searched for a gamma-ray signal from this object using the 35-hour observations taken by the MAGIC telescopes between November 2010 and January 2011. No significant VHE gamma-ray signal was detected. We computed the upper limits of the integrated gamma-ray flux at 95\\% confidence level of $9.4 \\times 10^{-12}$ cm$^{-2}$ s$^{-1}$ and $2.5 \\times 10^{-12}$ cm$^{-2}$ s$^{-1}$ above $100$ GeV and $200$ GeV, respectively. The differential upper limits in four energy bins in the range from $80$ GeV to $500$ GeV are also derived. The upper limits are consistent with the attenuation predicted by low-flux EBL models on the assumption of a simple power-law spectrum extrapolated from LAT data. ",
        "introduction": "\\label{sec:int}  Gamma-ray astronomy has rapidly grown in recent years thanks to the development of Imaging Atmospheric Cherenkov Telescopes (IACTs), which are able to observe gamma rays at very high energies (VHE; $> 30$ GeV), such as the High Energy gamma-ray Spectroscopic System \\citep[H.E.S.S.;][]{Aharonian2006AA457p899}, the Major Atmospheric Gamma-ray Imaging Cherenkov \\citep[MAGIC;][]{Aleksic2012APh35p435} telescopes, and the Very Energetic Radiation Imaging Telescope Array System \\citep[VERITAS;][]{Nepomuk2009ICRCpubID1408}. Space telescopes observing in the GeV energy band such as the Large Area Telescope (LAT) on board the {\\it Fermi} satellite \\citep{Atwood2009ApJ697p1071} have also substantially contributed to progress in the field.   VHE photons, with energies above a characteristic energy, are strongly attenuated through electron-positron pair creation in the extragalactic background light (EBL). This characteristic EBL absorption energy is defined by the photon energy at which the pair creation optical depth is unity. The characteristic energy depends on the source's redshift. The attenuation results in a much softer spectrum at Earth than that of the intrinsic spectrum of a source \\citep[e.g.,][]{Nikishov1962JETP14p393,Gould1966PRL16p252,Stecker1992ApJ390L49}. The characteristic energy is lower for sources with higher redshift. Thus, the detection of VHE gamma rays from distant sources is a challenge for IACTs. On the other hand, this absorption process brings a unique opportunity to study the EBL \\citep[e.g.,][]{Stecker1993ApJ415L71,Stanev1998ApJ494L159}, which is difficult to be measured directly due to strong Galactic foregrounds \\citep[e.g.,][]{Hauser2001ARAA39p249,Dwek2013APh43p112}. The EBL is mainly interpreted as stellar and dust emission integrated over the cosmic history. Therefore, VHE gamma-ray observations can be used to extract information on the imprinted cosmic stellar and galaxy evolution.   Observations of VHE gamma-ray sources with different redshifts allow us to investigate the EBL in different wavelengths and hence to examine underlying astrophysics and cosmology at different epochs. In particular, VHE gamma-ray observations of BL Lac objects at $z \\lesssim 0.2$ have constrained EBL models from optical to infrared wavelengths \\citep[e.g.,][]{Aharonian2006Nature440p1018,Aharonian2007AA475L9}. The most distant object ever detected in the VHE energy range, with firmly confirmed redshift, is the flat spectrum radio quasar (FSRQ) 3C 279 \\citep[$z = 0.536$;][]{Albert2008Sci320p1752}. This source has also been utilized to constrain EBL models on the assumption that the intrinsic power-law index in the VHE range is softer than $-1.5$. Recently, two sources with possibly higher redshifts have been detected, namely KUV 00311-1938 \\citep{Becherini2012AIPC1505p490} with $z \\geq 0.506$ \\citep{Pita2012AIPC1505p566} and PKS 1424+240 with $z \\geq 0.6035$ \\citep{Acciari2010ApJ708L100,Furniss2013ApJ768L31}, but they have not been used to estimate the EBL density due to the uncertainty in their redshift measurements. Moreover, different authors have recently claimed the detection of the EBL imprint by the statistical analyses of the spectra of blazars \\citep{Ackermann2012Sci338p1190,Abramowski2013A&A550A4,Dominguez2013ApJ770p77}. In particular, the analysis of \\citet{Ackermann2012Sci338p1190} covers the wide redshift range ($0.2 \\lesssim z \\leq 1.6$) using LAT data. \\citet{Abramowski2013A&A550A4} focus on lower redshift blazars detected by H.E.S.S. ($z < 0.2$). \\citet{Dominguez2013ApJ770p77} use the multi-wavelength data of the spectral energy distribution (SED) of blazars to estimate the cosmic gamma-ray horizon from the local Universe to $z \\sim 0.5$ by using an EBL-model-independent technique. Although the LAT has unveiled the EBL at high redshifts ($z \\sim 1$), the detection of sources with high redshifts in the VHE range allows us to test different (longer than in the LAT energy range) wavelengths of the EBL which have not yet been investigated. The MAGIC telescopes have achieved an analysis energy threshold below 100 GeV, and therefore are an ideally suited instrument among existing IACTs to study high-redshift sources.   In this paper we report the observational results of the distant quasar 4C +55.17 ($z = 0.896$) by the MAGIC telescopes. This source was recognized as a promising high-redshift candidate for VHE emission from the first LAT AGN Catalog \\citep[1LAC;][]{Abdo2010ApJ715p429}. Moreover, 4C + 55.17 was also identified as a good source candidate for VHE detection by the {\\it Fermi} LAT collaboration in a dedicated high-energy data analysis (private communication, September 2010), and by \\citet{Neronov2011AA529A59}. The spectrum in the LAT energy range is hard and consistent with a power-law $dN / dE \\propto E^{-\\alpha}$ with $\\alpha = 2.05 \\pm 0.03$ in the range from $100$ MeV to $100$ GeV, not showing significant flux variability in the 1LAC \\citep{Abdo2010ApJ715p429}. The spectral hardness of gamma-ray sources is particularly important to study the EBL because of the limited sensitivities of IACTs; harder sources can provide more photons with energies above the characteristic EBL absorption energy, allowing for studies with better precision. Even with a spectral break at the energy of several GeV recently reported \\citep{Ackermann2011ApJ743p171,McConville2011ApJ738p148,Neronov2012arXiv1207.1962} this source is still a good candidate for a distant VHE gamma-ray emitter.   This source is a radio-loud active galaxy with the firmly confirmed redshift of $z = 0.896$ estimated from broad optical emission lines in its spectrum \\citep{Wills1995ApJ447p139,Adelman-McCarthy2008ApJS175p297}. It had been categorized as a FSRQ, but interestingly the classification has recently been questioned due to its morphological and spectroscopic properties which are distinct from those of other FSRQs \\citep{Rossetti2005AA434p449}. FSRQs are characterized by a centrally concentrated radio core with high brightness temperature and highly variable flux. However, the brightness temperature of this source is $2 \\times 10^{8}$ K at 5 GHz \\citep{Taylor2007ApJ671p1355}, which is much lower than that of the other known quasar-hosted gamma-ray blazars \\citep{McConville2011ApJ738p148}, and a light curve in the LAT energy range is consistent with steady emission \\citep{Furniss2013arXiv1303.1103}. Note that only strong variability can be significantly detected due to the limited number of detected photons by the LAT. Radio morphology has indicated that 4C +55.17 belongs to a family of young radio galaxies, that is, compact symmetric objects \\citep[see][for a review]{O'Dea:1998mc}, which are a smaller version of classical Fanaroff-Riley (FR) II radio galaxies \\citep{Rossetti2005AA434p449}. The low flux variability of this source also supports this indication. Importantly, this type of compact objects is a promising source candidate of ultra-high-energy cosmic rays \\citep{Takami2011APh34p749} similarly to the lobes of FR II radio galaxies \\citep[e.g.,][]{Biermann1987ApJ322p643,Takahara1990PTP83p1071,Rachen1993AA272p161}. \\citet{McConville2011ApJ738p148} model the multi-wavelength SED of this source by using both compact symmetric object and FSRQ interpretations. The authors suggest that infrared and hard X-ray observations can help to distinguish between the two models, but VHE gamma rays also provide a clue for the distinction despite the significant EBL absorption. Also, a hadronic interpretation predicts a hard intrinsic spectrum in the VHE range \\citep{Kino2011MNRAS412L20}. Therefore, observations of this source in the VHE range can provide an important clue to its nature and radiation mechanism.   This paper is laid out as follows. We describe the details of the MAGIC observations and data analysis in Section \\ref{sec:obs}. Then, results are shown in Section \\ref{sec:res}. Implications from the results are discussed in Section \\ref{sec:dis} and a summary is presented in Section \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  The MAGIC telescopes observed 4C +55.17 for $35$ hours ($27.73$ hours of effective time with good quality data) from November 2010 to January 2011 in a wobble mode. No significant gamma-ray signal was found above $100$ GeV. Instead, both integral and differential upper limits of gamma-ray flux were derived. The integral upper limits are $9.4 \\times 10^{-12}$ cm$^{-2}$ s$^{-1}$ above $100$ GeV and $2.5 \\times 10^{-12}$ cm$^{-2}$ s$^{-1}$ above $200$ GeV. The differential upper limits are tabulated in Table \\ref{tab:difful}. The derived limits are close to and consistent with the power-law spectrum extended from the LAT energy range attenuated by low-flux EBL models.   \\citet{McConville2011ApJ738p148} claimed that a 50-hour observation with MAGIC would suffice to detect this source. However, the strict upper limits from $\\sim 30$ hours of observation do not confirm this prediction, due to the medium zenith angle of the observations ($\\gtrsim 27^{\\circ}$), which increases the analysis threshold and suppress the performance at lowest energies \\citep{Aleksic2012APh35p435}. In fact, according to our estimates, MAGIC would require $\\sim 100$ hours in order to have a $5 \\sigma$ detection from 4C +55.17 under the optimistic assumption of the power-law extension adopted in this study and relatively optically-thin EBL models such as \\citet{Franceschini2008AA487p837} and \\citet{Dominguez2011MNRAS410p2556}. After the observation of 4C +55.17, the camera of the MAGIC I telescope was upgraded, which improved the sensitivity for the lowest energies \\citep{Sitarek2013arXiv1308.0141}. A new trigger system for stereoscopic observations, so-called sum-trigger \\citep{Haefner2011ICRC9p246}, will improve the sensitivity below $100$ GeV significantly. These may help detect this source in the VHE range. Detailed studies on its spectrum beyond detection would require a sensitivity several times better than current ground-based gamma-ray instruments, and hence it may be a task for the next generation of instruments such as the Cherenkov Telescope Array \\citep[CTA;][]{Actis2011ExA32p193,Acharya2013APh43p3}."
    },
    "1402/1402.1496_arXiv.txt": {
        "abstract": "Strong gravitational lenses can be used to detect low mass subhalos, based on deviations in image fluxes and positions from what can be achieved with a smooth mass distribution. So far, this method has been limited by the small number of (radio-loud, microlensing free) systems which can be analysed for the presence of substructure. Using the gravitational lens B1422+231, we demonstrate that adaptive optics integral field spectroscopy can also be used to detect dark substructures.  We analyse data obtained with OSIRIS on the Keck I Telescope, using a Bayesian method that accounts for uncertainties relating to the point spread function and image positions in the separate exposures. The narrow-line [OIII] fluxes measured for the lensed images are consistent with those measured in the radio, and show a significant deviation from what would be expected in a smooth mass distribution, consistent with the presence of a perturbing low mass halo. Detailed lens modelling shows that image fluxes and positions are fit significantly better when the lens is modelled as a system containing a single perturbing subhalo in addition to the main halo, rather than by the main halo on its own, indicating the significant detection of substructure. The inferred mass of the subhalo depends on the subhalo mass density profile: the 68 \\% confidence interval for the perturber mass within 600 pc are: 8.2$^{+0.6}_{-0.8}$ , 8.2$^{+0.6}_{-1}$ and 7.6$\\pm0.3$ $\\log_{10}$[M$_{\\rm{sub}}$/M$_\\odot$] respectively for a singular isothermal sphere, a pseudo-Jaffe, and an NFW mass profile. This method can extend the study of flux ratio anomalies to virtually all quadruply imaged quasars, and therefore offers great potential to improve the determination of the subhalo mass function in the near future. ",
        "introduction": "\\label{sec:intro}  A fundamental prediction of the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model is that a Milky Way mass halo should be surrounded by thousands of subhalos. However, only about twenty are observed \\citep[e.g.][]{Klypin++99, Moore++1999, Strigari++07, Toll++08, Koposov++08}. Significant work in the past decade has gone into understanding this failure in a theory that has otherwise been successful at reproducing the observed power spectrum of structure in the Universe over an enormous range of distance, mass and time scales \\citep[e.g.][]{Planck++13,Komatsu++11}.  If CDM is correct, there must be a large number of subhalos which do not retain enough gas \\citep{Papastergis++11} or form enough stars to be detected. A variety of baryonic processes are suspected to lead to dark subhalos. In fact, star formation in satellite galaxies has been suggested to be quenched by a variety of internal and external process such as UV heating during reionization, tidal and ram-pressure stripping by the central galaxy, supernova feedback, and stellar winds \\citep[e.g.][]{Lu++12, Menci++12, Guo++11,Bullock++2000,Benson++2002, Somerville++02, Kravtsov++04, Kaufmann++08, Maccio++10,Springel++10,Zolotov++12}. Significant uncertainty as to the relative importance of various star formation processes remains, with numerous models reproducing the luminosity function of Milky Way satellites while failing to match the satellite luminosity function in different host mass and redshift regimes \\citep{Nierenberg++13}. An alternative explanation for the lack of satellites is that dark matter subhalos are not as numerous as predicted by $\\Lambda$CDM, as in the case of warm dark matter models \\citep{Nierenberg++13}.  There are also discrepancies between the subhalo mass profiles predicted by dark matter only CDM simulations and what is observed in Local Group dwarf galaxies. \\citet{Boylan-Kolchin++11, Boylan-Kolchin++12}, pointed out that the most luminous Milky Way satellites reside in halos which have much lower circular velocities than what is predicted by simulations. Baryonic processes could account for a flattening of the central density profile in the more massive satellite galaxies \\citep{Zolotov++12}. Alternatively, self-interacting cold dark matter (SIDM) has been proposed to explain the flattening \\citep[e.g.][]{Rocha++13}.   The only way to fundamentally distinguish between these various dark matter models is with a direct measurement of the subhalo mass function which does not rely on an understanding of the baryonic physics in the subhalos. In the Local Group, new work has shown the possibility that dark subhalos may be detected via their interactions with tidal streams \\citep{Carlberg++12}, or the HI disk of the Milky Way \\citep{Chakrabarti++09}. However, these methods are currently limited to low redshifts, and rely on detailed modelling of baryonic structures.  Outside of the low redshift Universe, gravitational lensing can be used to measure the subhalo mass function at a range of redshifts, without requiring detailed modelling of baryonic physics. Weak gravitational lensing can be used to measure the masses of massive ($\\sim 10^{11.5}$M$\\odot$) satellites of galaxy groups \\citep{Li++14}, while strong gravitational lensing is sensitive to lower mass subhalos. In strong gravitational lensing, a background source is multiply imaged by a foreground deflector, with the image time-delays, positions, and magnifications depending sensitively on the gravitational potential of the deflector, and its first and second derivatives respectively. Precise astrometry and photometry of lensed images have been shown to be powerful ways to measure the mass of low-mass subhalos in proximity to lensed images \\citep[e.g.][]{Mao++98,M+M01,Dalal++02,Koo05, Ama++06,McKean++07, K+M09, MacLeod++09,Treu++10,Vegetti++10a,Vegetti++12,Fadely++12}.  There are two main requirements for a strong gravitational lens to be suitable for the detection of substructure. First, the background source must be sufficiently large that it is not significantly magnified by stars in the plane of the lens galaxy. For typical lens configurations, stars have Einstein radii of order microarcseconds, thus the background source must be significantly larger than microarcseconds in apparent size. Secondly, the lensed images must contain enough information to constrain the lensing `macromodel', i.e., the smooth mass distribution of the main deflector.  Traditionally, these restrictions have limited the study of substructure to the seven known radio-loud quasar sources which are quadruply imaged (quad).  In a seminal work, \\citet{Dalal++02} demonstrated that these systems contained a fraction of mass in substructure which was broadly consistent with predictions from CDM simulations, albeit with large uncertainties. Progress with this technique has been limited by the small number of known suitable systems.  An alternative method of using gravitational lensing to study substructure is gravitational imaging, proposed by \\citet{Koo05}, in which subhalos are detected via astrometric perturbations of lensed galaxies. This method has been successfully applied to HST and AO imaging \\citet{Vegetti++12, Vegetti++10b}, and simulations indicate it will also work for ALMA data \\citep{Hezaveh++13}. Recently \\citet{Vegetti++14} analysed a sample of 11 lenses, finding a fraction of substructure consistent with numerical simulations and results from analyses of four image quasar lenses. The mass sensitivity of this method is determined by the signal to noise ratio and resolution of imaging, as well as the intrinsic morphology of the lensed galaxy. Current measurements with Keck AO are sensitive to substructure with masses of $\\sim 10^{7.5}$M$_\\odot$ and larger, while measurements with the Next Generation Adaptive Optics on Keck, the Thirty Meter Telescope, Very Long Baseline Interferomety may allow the sensitivity of this method to reach masses of $\\sim 10^6$M$_\\odot$.   Strongly lensed \\emph{narrow-line} quasar emission provides an alternative means of measuring the subhalo mass function \\citep{Moustakas++03}. Quasar narrow-line emission at low redshift is observed to be typically extended over tens to hundreds of parsecs depending on the source luminosity \\citep{Bennert++02}, which corresponds to milliarcseconds for typical source redshifts, well above the microlensing scale. A benefit to studying narrow-line emission is that many more quasars have detectable narrow-line emission as opposed to radio emission. This makes this method ideal for measuring the subhalo mass function in the thousands of quad quasar lenses which are expected to be discovered based on their optical properties in ongoing and future surveys including PANSTARRS, DES and LSST \\citep{Oguri++10}.  While promising, spatially resolved spectroscopy of narrow line emission of quasar lenses has proven difficult to attain. \\citet{Sluse++12a} and \\citet{Guerras++13} have measured lensed image spectra for pairs of quasar images, and for wide image separation double lenses, and observed differential lensing between the broad and continuum emission, which is evidence for microlensing. In order to study the lensing signal of substructure as discussed before, it is necessary to measure the spectra of images individually. \\citet{Metcalf++04} and \\citet{Keeton++06}, achieved sufficient spatial resolution in order to measure the broad emission in individual quad quasar images, but neither of these studies detected narrow emission with sufficient signal to noise ratio.  In a pioneering work, \\citet{Sugai++07} used integral field unit (IFU) optical spectroscopy in order to measure narrow [OIII] fluxes in the gravitational lens RXJ1131. The relatively wide separation between the lensed images ($\\sim 1\\farcs$0) made it possible to spatially resolve the lensed images, and place an upper limit on the mass of perturbing substructure. Interestingly, they found that unlike the continuum and broad fluxes, the [OIII] image fluxes did not deviate significantly from a smooth model prediction without substructure. In order to apply this method to a larger sample of systems and to probe lenses with configurations with smaller image separations (such as fold lenses) it is essential to obtain higher spatial resolution via HST \\citep{Keeton++06}, or adaptive optics from the ground. In this work, we demonstrate for the first time that subhalos can be detected using strongly lensed narrow-line quasar emission of a high redshift (z=3.6) quasar by combining adaptive optics with integral field spectroscopy in order to obtain spatially resolved spectra of individual lensed images.   In this paper we present a measurement of narrow-line lensing with unprecedented accuracy using the OH Suppressing Infra-Red Imaging Spectrograph \\citep[OSIRIS][]{Larkin++06} with laser guide star adaptive optics (LGS-AO) \\citep{Wizinowich++06} at Keck.  We choose as initial case to study the famous system B1422+231 \\citep{Patnaik++92}. Previous extensive studies across many bands of photometry (including radio) enable us to compare our measurement with more traditional methods.  The structure of the paper is as follows: In \\S \\ref{sec:data} we present our observations. In \\S  \\ref{sec:imExtract} we describe our Bayesian method for combining multiple exposures to get sub-pixel sampling of the images and extract spectra optimally. In \\S \\ref{sec:spectra} we present the extracted spectra and show fits to the broad and narrow line components in each of the lensed images, and the inferred fluxes from each. In \\S \\ref{sec:lensmodels} we discuss the gravitational lens models we use to fit the observed image fluxes and positions. In \\S \\ref{sec:results} we present results from a gravitational lens model assuming a single perturbing subhalo. In \\S \\ref{sec:discussion} we discuss future prospects for this method. Finally in \\S \\ref{sec:summary}, we summarise our results. Throughout this paper, we assume a flat $\\Lambda$CDM cosmology with $h=0.7$ and $\\Omega_{\\rm m}=0.3$.  All magnitudes are given in the AB system \\citep{Oke++1974} unless otherwise stated. ",
        "conclusions": "\\label{sec:discussion}  Our primary goal in this work is to demonstrate that strongly lensed narrow-line emission provides an alternative to radio emission in lensed quasars for detecting substructure, yielding lens model constraints with consistent results and comparable accuracy. In this section we discuss previous results for this system, how the analysis can be improved by considering additional systematic uncertainties, as well as future prospects for the method  \\subsection{Comparison with previous work} B1422+231 is a bright, radio loud system which enabled one of the earliest detections of dark substructure at cosmological distances \\citep[e.g.][]{Mao++98,Keeton++01, Bradac++05}. These studies found that the observed radio fluxes could be best explained by a perturbing substructure with characteristic Einstein radius of a few mas. Although there was some question about how macromodel assumptions might affect the inference of substructure \\citep{Kawano++04}, \\citet{Kochanek++04} demonstrated that smooth macromodels that are consistent with observations of galaxies using weak lensing and galaxy morphology do not provide a good fit to the observed image positions and flux ratios. Our analysis based on lensed narrow-line emission yields consistent results to these studies.  \\subsection{Limitations of our analysis} We have explored how our data can be used to constrain a mass model with a main lens in external shear and one perturbing subhalo. As we discussed in Section \\ref{sec:perturbed}, this model makes several important simplifying assumptions.  First, for simplicity we adopt a macro-model composed by a singular isothermal ellipsoid with external shear. This is in general found to be a good model for early-type galaxies both in terms of radial profile \\cite{Koo++06,Koo++09} and angular structure \\citep{Yoo++05,Yoo++06}. However, allowing for more flexibility in the model could inflate the uncertainties on the inferred substructure, or reduce its significance, even though in the case of B1422+231 it is hard to find physically plausible models that fit the data without a localized substructure \\citep{E+W03}.  Second, as with previous works, we explore the effects of a single perturbing subhalo \\citep{MacLeod++09,McKean++07,MacLeod++13}. When studying the effects of subhalos on lensed compact image fluxes and positions, full populations of subhalos are computationally expensive and not well constrained by a single lensing system \\citep{Fadely++12}. In the future, when large samples will be available, it will be important to carry out a systematic analysis of the entire population of subhalos using hierarchical modelling to infer at once the mass function and spatial distribution of the subhalos.  Third, similarly to other works of this type, we have assumed that the perturbing substructure is in the plane of lens galaxy in order to estimate the perturber mass \\citep[e.g.][]{Dalal++02, MacLeod++09,Fadely++12,Vegetti++10a,Vegetti++12,Vegetti++14}. \\citet{Xu++12} showed that line of sight structure may contribute significantly to image magnifications, although the observed frequency of radio flux ratio anomalies can be explained with in situ perturbers predicted by CDM \\citep{Xu++14, Metcalf++12}. This has important implications for the inferred subhalo mass function as all masses are inferred assuming the redshift of the perturber is known. The lensing effect varies with mass if the perturber is not within the plane of the lens galaxy \\citep{Xu++12,McCully++2014}.  We have also demonstrated the effects of varying the assumed mass profile of the perturbing subhalo, which significantly change the posterior probability distribution for the perturber mass and position. It would be interesting in a future work to consider a range of subhalo mass profiles which are drawn directly from simulations, so that differential tidal stripping as a function of three dimensional distance from the lens halo centre, and scatter in the subhalo concentration mass relation can be taken into account for instance.  \\subsection{Future Prospects} Many more quasars have significant narrow-line emission than radio emission, thus this method can be extended to study substructure in a larger sample of systems. There is significant future potential for this method as ongoing and planned optical surveys such as DES, LSST, PANSTARRS and GAIA are expected to find thousands more lensed quasars based on their optical rather than radio properties \\citep{Oguri++10}. The next generation of adaptive optics systems (NGAO), \\citep{Max++08}, and telescopes such as the Thirty Meter Telescope and the European Extremely Large Telescope will make it possible to measure narrow-line lensing rapidly and with high precision for those large samples of objects. In the nearer future, upgrades to the Keck AO system will enable PSF telemetry which will thereby make it possible to obtain better constraints on the source size, and thus place stronger lower limits on perturber masses.           Significant information is also contained in deep imaging of the lensed systems, which can sometimes reveal luminous substructure \\citep{McKean++07, MacLeod++09}. If a luminous satellite is detected, it can be used to break the degeneracy between the subhalo mass and position, as well as providing a measurement of the subhalo mass to light ratio.    We used OSIRIS at Keck with adaptive optics in order to obtain spatially resolved spectroscopy of the gravitational lens B1422+231. We developed a new pipeline which enabled us to robustly extract lensed image spectra, by first inferring the PSF and image positions in each exposure. Using this information, we detected a significant deviation from a smooth gravitational lens model, which we used to infer the presence of a perturbing low mass subhalo. Our main results are summarised below:  \\begin{enumerate} \\item Using our data reduction pipeline, we were able to measure image positions to 5 mas accuracy and integrated line flux ratios with $\\sim$3\\% uncertainties.  \\item The narrow-line flux ratios are consistent with radio measurements, and deviate significantly from the smooth model predictions. The broad-line flux ratios are consistent with the narrow line measurements.  Our measurement of the continuum is significantly offset, showing a much stronger deviation from the smooth model, possibly indicative of microlensing in addition to millilensing.  \\item % Based on the assumption that the deviation of observed narrow-line fluxes from the smooth model is due to a single perturbing subhalo, we infer the mass and position of the perturbing subhalo for three different mass profiles. The mass profile significantly affects possible perturber positions and masses, with an NFW subhalo restricted to the smallest region of parameter space by its shallower mass profile which has a relatively weaker lensing signal. The 68\\% confidence intervals for the logarithmic perturber mass within 600 pc are:  8.2$^{+0.6}_{-0.8}$, 8.2$^{+0.6}_{-1}$ and 7.6$\\pm 0.3$ $\\log_{10}$[M$_{\\rm{sub}}$/M$_\\odot$], for a singular isothermal sphere, pseudo-Jaffe, and NFW mass profile respectively.  The mass within 300 pc is 7.8$^{+0.6}_{-0.7}$ and 8.0$^{+0.6}_{-0.8}$, 7.2$^{+0.2}_{-0.2}$$\\log_{10}$[M$_{\\rm{sub}}$/M$_\\odot$]  \\end{enumerate}"
    },
    "1402/1402.6755_arXiv.txt": {
        "abstract": "We have identified spectral features in the late-time X-ray afterglow of the unusually long, slow-decaying \\grb using \\nustarn, \\swiftn-XRT, and \\chandran.   A spectral component in addition to an absorbed power-law is required at $>4\\sigma$ significance, and its spectral shape varies between two observation epochs at $2\\times10^5$ and $10^6$ seconds after the burst. Several models can fit this additional component, each with very different physical implications.  A broad, resolved Gaussian absorption feature of several keV width improves the fit, but it is poorly constrained in the second epoch.  An additive black body or second power-law component provide better fits.  Both are challenging to interpret: the blackbody radius is near the scale of a compact remnant ($10^8$\\,cm), while the second powerlaw component requires an unobserved high-energy cutoff in order to be consistent with the non-detection by \\textit{Fermi}-LAT. ",
        "introduction": "Recent work has identified several ``ultra-long'' gamma-ray bursts (GRBs) with properties distinct from normal long GRBs \\citep[and references therein]{tmp_Levan:13:ULGRBs}.  These events have initial bursting phases lasting thousands of seconds in gamma-rays and show long-lived, highly variable X-ray afterglows.  It is currently unclear whether these bursts are simply extreme examples of the long GRB class, as suggested by \\citet{tmp_Zhang:13:ULGRBs}; if they are related to the even longer candidate relativistic Tidal Disruption Events (TDEs) \\swift J1644+57 \\citep{Bloom:11:SwiftJ1644TDF,Levan:11:SwiftJ1644TDF, Burrows:11:SwiftJ1644TDF} and \\swift J2058+05 \\citep{Cenko:12:TDESwiftJ2058}; or if they represent a new subclass of transient, perhaps with large-radius progenitors \\citep{Woosley:12:LongCollapsars,Gendre:13:ULGRB111209A,Nakauchi:13:BSGULGRB}.  The bright, nearby \\grb is similar to previously reported ultra-long GRBs and, with the launch of \\nustarn, provides an opportunity to observe the X-ray spectrum at high sensitivity over a broad energy band. Here we report time-varying spectral features in the late-time X-ray afterglow of \\grb that were initially discovered by \\nustar and confirmed in a second epoch by \\nustar and \\chandran. Our detections are at higher energies and significantly later times than previously reported afterglow features.  Before the era of routine afterglow observations with \\swiftn-XRT, several authors claimed detection of lines in GRB X-ray afterglows on top of otherwise smooth power-law spectra \\citep[e.g.,][]{Piro:2000:AGLine,Amati:00:AGAbsLine,Reeves:02:AGLine}. Most reports were of emission lines at relatively low signal-to-noise ratio, and there was substantial controversy over the methods used to assess line significance \\citep{Protassov:02:LRTest, Sako:05:AGLineReanalysis}. Since the advent of \\swiftn, no firm afterglow line detections have been reported despite its greater sensitivity and systematic followup, calling previous reports into question \\citep[for a review, see][]{Hurkett:08:XRTLineSearches}.  However, statistically significant blackbody components have been reported in the early-time ($t \\lesssim 10^3$\\,s) afterglow spectra of several bursts observed by \\swiftn-XRT \\citep[][and references therein]{Starling:12:AGBBody}. The inferred rest-frame temperatures are typically a few tenths of a keV, the inferred radii are $\\sim10^{12}$\\,cm, and the blackbody component provides 10--50\\% of the 0.3--10\\,keV flux. The first detections were in low-luminosity, SN-associated GRBs, leading to suggestions that the emission was due to shock breakout from the SN \\citep[e.g.,][]{Campana:06:ThermalAG}.  % Systematic searches have found thermal components in early afterglows of classical GRBs as well \\citep{Sparre:12:AGBBSearch,Friis:13:AGBBSearch}, giving credence to alternative interpretations including late-time emission from a prompt photosphere \\citep{Friis:13:AGBBSearch} or emission from a cocoon around the jet \\citep{Suzuki:13:ThermalCocoon,Nakauchi:13:BSGULGRB}.  Of particular relevance are reports of additional components in the afterglows of other ultra-long GRBs. The ``Christmas Day Burst'' GRB\\,101225A showed evidence of two separate blackbody components, a 1\\,keV X-ray black body with radius $2\\times10^{11}$\\,cm observed 6\\,ksec after the burst and a UVOIR black body with radius 2--7$\\times10^{14}$\\,cm which cooled over 18\\,days \\citep{Thone:11:ChristmasDayBurst}.  In GRB\\,111209A, \\citet{Stratta:13:ULGRB111209A} reported the \\textit{XMM} detection of a second, hard power-law component ($\\Gamma \\sim 0$) during the steep decay phase $\\sim$70\\,ksec after the burst. ",
        "conclusions": "\\label{sec:discussion}  Our late-time afterglow observations of \\grb require an additional spectral component at high significance.  Several alternative models provide acceptable fits to the data. These spectral features are detected more than 1\\,Msec after the burst, much later than any components previously reported in X-ray afterglows, probing a largely unexplored phase of afterglow evolution. Several unique features of \\grb make it possible to detect these late-time features for the first time.  The unusually bright afterglow enables high-quality spectral fits, and \\nustar has excellent sensitivity at the relevant energies and can constrain the continuum above 10 keV.  Moreover, the primary power law is unusually soft, so the high-energy component is not swamped.  It is not yet clear whether this emission is related to progenitor physics unique to this unusual, ultra-long burst; \\nustar observations of the bright ``canonical'' long GRB\\,130427A were consistent with emission by a single spectral component \\citep{Kouveliotou:13:GRB130427A}.  Future observations of bright afterglows will be needed to determine the prevalence of these late-time spectral components and identify the relevant emission mechanism."
    },
    "1402/1402.4616_arXiv.txt": {
        "abstract": "We present a review of modern optical techniques being used and developed for the field of gravitational wave detection. We describe the current state-of-the-art of gravitational waves detector technologies with regard to optical layouts, suspensions and test masses. We discuss the dominant sources and noise in each of these subsystems and the developments that will help mitigate them for future generations of detectors. We very briefly summarise some of the novel astrophysics that will be possible with these upgraded detectors. ",
        "introduction": "\\label{sec:gws} Over the last century our ability to view the Universe has undergone many revolutions. Observational astronomy has progressed from a time when the only tool was ground-based visible light photometry to an era of ground-and-space-based multi-wavelength photometry and spectroscopy. These advances have led to countless new discoveries and greatly increased our knowledge of the Universe, but have only been possible through major technological and scientific leaps in instrumentation. Sometimes this has been through serendipitous applications from different fields, but often it has been the astronomical need that has pushed the technological boundaries.  Gravitational waves (also called gravitational radiation) offer another new window on the Universe moving beyond the electromagnetic spectrum. They are a direct prediction of Einstein's General Theory of Relativity (GR), although their main properties can be inferred using Newtonian gravity and the principal of special relativity that information (including the influence of gravity) cannot propagate faster than the speed of light \\cite{Schutz:1984}. As they travel through space gravitational waves manifest as a transverse varying tidal strain, which would have the effect of altering the proper distance between freely falling objects. Gravitational waves are quadrupolar in nature and have two polarisation states (called `plus' [$+$] and `cross' [$\\times$]) rotated by 45$^\\circ$ with respect to each other (see Figure~\\ref{fig:ringofparticles}). These states have orthogonal polarisation.  \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\textwidth]{figs/ringofparticles.eps} \\caption{The effect of a passing gravitational wave on a ring of freely falling test particles. The top and bottom panels show the effect of the `plus' and `cross' polarisation modes respectively.} \\label{fig:ringofparticles} \\end{center} \\end{figure}  Gravity is the weakest of the fundamental forces and correspondingly the effect of \\gws is tiny. Their strength is often measured as a strain in space, where for our ring of particles in Figure~\\ref{fig:ringofparticles} a maximum displacement of $\\Delta l$, given an un-displaced radius of $l$, means a strain of $h = \\Delta l/l$. For a source at a distance $r$ the approximate strain amplitude is given by \\begin{equation} h \\approx \\frac{2GM}{c^2r}\\left(\\frac{v}{c}\\right)^2, \\end{equation} where $v$ is the non-spherical component of the source's motion (spherically symmetric motions, e.g.\\ a rotating perfect sphere will not radiate as they lack a time-varying quadrupole moment) and $M$ is the mass involved in that motion. It can be seen that even for nearby sources (e.g.\\ $1\\,{\\rm kiloparsec} = 3.09\\ee{19}$\\,m) with the mass of the Sun (1$M_{\\odot} = 1.99\\ee{30}$\\,kg) moving at near-relativistic velocities ($0.01c = 3.00\\ee{6}$\\,m/s) the strain at Earth will be $h \\sim 3\\ee{-20}$, equivalent to a displacement between the Sun and the Earth of $\\sim 5$\\,nm.  As of yet no \\gws have been directly detected. However, their existence has long been indirectly inferred from the excellent agreement between theory and observation of the orbital decay of neutron star binary systems (e.g.\\ \\cite{Taylor:1982}).  \\subsection{Astrophysical sources of gravitational waves} Any non-axisymmetric acceleration of mass will emit \\gws, but as seen above the emission of strong (and potentially observable) \\gws requires very large masses and relativistic motion. To fulfil both these requirements the sources must be compact and dense. The archetypal source is the coalescence of a pair of compact stars (e.g.\\ neutron stars or black holes) in a tight binary system. Other sources include: individual rotating neutron stars with sustained distortions on them \\cite{Zimmermann:1979,Bonazzola:1996}; vibrating neutron stars \\cite{Andersson:1998} and black holes \\cite{Teukolsky:1973}; the collapsing core of a star during a supernova \\cite{Ott:2009} or $\\gamma$-ray burst \\cite{Corsi:2012}; or, more exotically, inflation following the Big Bang \\cite{Rubakov:1982} and cosmic strings \\cite{Vilenkin:1981,Damour:2000}. An excellent review of \\gw sources and the science that can be learned from them is presented in \\cite{Sathyaprakash:2009}.  \\subsection{The detection of gravitational waves} The basic principle behind any \\gw detector is to measure the relative displacement of freely falling bodies. The first attempts to construct a detector to observe them were made by Joseph Weber in the 1960s \\cite{Weber:1969}. He used resonant mass, or ``bar'', detectors consisting of a suspended aluminium cylinder encircled with piezoelectric sensors, and designed to resonate within a narrow frequency band when excited by a passing \\gw. Modern bar detectors utilised improved seismic isolation and operated at cryogenic temperatures ($\\leq6$\\,K) to reduce the thermal noise. Detectors including ALLEGRO (Louisiana, US), EXPLORER (CERN), NAUTILUS (Frascati, Italy) and NIOBE (Perth, Australia) have been used to set interesting limits on gravitational wave emission \\cite{Ronga:2006}. Following on from these inherently narrow-band ``bar'' detectors came the broad-band interferometric \\gw detectors \\cite{Moss:1971, Weiss:1972, Drever:1977, Billing:1979}, which were pioneered by Rai Weiss at MIT \\cite{Weiss:1972}. The aim of these was to use interferometry to measure the displacement of two suspended mirror-coated test masses at the end of each arm. The techniques required to allow these detectors to achieve the extraordinary accuracy necessary to detect a \\gw will be detailed within this review. One thing that can be seen straight away from the previous definition of the strain, $h$, is that larger displacements, or smaller strains, can be measured if the separation of the test masses is large. Therefore a prerequisite of current and future generations of interferometric \\gw detectors is that the arm lengths are long, generally on the few kilometres scale.  \\subsection{Searches for gravitational waves}\\label{sec:gwsearches} The most recent efforts to detect \\gws have come from the US Laser Interferometric Gravitational-wave Observatory (LIGO) \\cite{Abramovici:1992, Abbott:2009b}, the French-Italian Virgo \\cite{Bradaschia:1990, Acernese:2006} and the British-German GEO600 \\cite{Willke:2007} projects. LIGO operates two sites (one in Livingston, Louisiana called LLO and one in Hanford, Washington called LHO) each housing an interferometer with 4\\,km arms and Fabry-Perot cavities (previously LHO also housed a 2\\,km interferometer). Virgo is a 3\\,km interferometer with Fabry-Perot cavities in Cascina, Italy. GEO600 (or latterly GEO-HF \\cite{Willke:2006, Lueck_HF}), located near Hannover, Germany, has 1.2\\,km arms, but in a folded configuration and with no Fabry-Perot cavities.  Over the last decade these detectors have produced unprecedented limits on \\gw emission from a wide variety of sources. The sensitivity of these detectors has meant that they would be able to observe neutron star binary coalescences out to distances of a few tens of Mpc \\cite{Abadie:2010} (a review of expected source rates can be found in \\cite{Abadie:2010b}). Some highlights of the astrophysics achieved include: surpassing the spin-down limit (assuming the star's observed slow down is entirely due to loss of energy through \\gws) for the Crab \\cite{Abbott:2008, Abbott:2010} and Vela \\cite{Abadie:2011} pulsars; excluding binary coalescences as progenitors of two nearby $\\gamma$-ray bursts \\cite{Abbott:2008b, Abadie:2012}; and, limiting the stochastic \\gw background to have an energy density less than $7\\ee{-6}$ of the closure density of the Universe \\cite{Abbott:2009}. A review of many more of the searches can be found in \\cite{Pitkin:2011}.  The LIGO and Virgo detectors are currently undergoing upgrades to their ``advanced'', or second generation, configurations: Advanced LIGO (aLIGO) \\cite{Harry:2010} and Advanced Virgo (AdV) \\cite{Acernese:2009}. An estimate of the commissioning and observation schedule for these advanced detectors is discussed in \\cite{Aasi:2013}. Another second generation interferometric detector, which will be located underground and utilise cryogenic cooling of its mirror, is the Japanese KAGRA \\cite{Somiya:2011} (formerly the Large-Scale Cryogenic Gravitational-wave Telescope, or LCGT \\cite{Kuroda:2010}). These detectors will aim to increase sensitivity over the initial generation of detectors by a factor of ten, giving a thousandfold increase the observable volume, in addition to lowering the operating frequency down to $\\simeq 10$\\,Hz. Theoretical predictions suggest \\cite{Abadie:2010b} that the likely number of observable neutron stars coalescences for these detector, once at design sensitivity, could be 40 per year, with the most optimistic estimates suggesting over one per day.  One fundamental limit to the low-frequency ($\\lesssim 1$\\,Hz) sensitivity of ground-based detectors is Newtonian noise, so to observe sources below this (e.g.\\ supermassive black hole binaries \\cite{Haehnelt:1994} and extreme mass ratio black hole inspirals \\cite{Gair:2004}) requires detectors in space like the proposed evolved Laser Interferometer Space Antenna (eLISA) \\cite{Seoane:2013, Jennrich:2009}, or similar (e.g.\\ \\cite{Kawamura:2011, Harry:2006, Ni:2009}). Detection of nanoHz \\gws from the most massive supermassive black hole systems is also possible via use of high-precision pulsar timing \\cite{Hobbs:2010}, in which pulsars themselves are used as the test masses for a galactic-scale detector. Even lower frequency \\gws could be imprinted in the Cosmic Microwave Background (CMB) and be detectable through the polarisation state of the microwave photons. Recent observations by the BICEP2 Collaboration show the first possible evidence for such relic gravitational waves in the CMB \\cite{Ade:2014}, although confirmation with other independent observations would strengthen the detection (e.g.\\ with the Planck satellite \\cite{Ade:2013, Zhao:2010}). We will not discuss any of these techniques any further in this review. ",
        "conclusions": ""
    },
    "1402/1402.0035_arXiv.txt": {
        "abstract": "The observable macroscopic properties of relativistic stars (whose equations of state are known) can be predicted by solving the stellar structure equations that follow from Einstein's equation. For neutron stars, however, our knowledge of the equation of state is poor, so the direct stellar structure problem can not be solved without modeling the highest density part of the equation of state in some way.  This talk will describe recent work on developing a model independent approach to determining the high-density neutron-star equation of state by solving an inverse stellar structure problem.  This method uses the fact that Einstein's equation provides a deterministic relationship between the equation of state and the macroscopic observables of the stars which are composed of that material.  This talk illustrates how this method will be able to determine the high-density part of the neutron-star equation of state with few percent accuracy when high quality measurements of the masses and radii of just two or three neutron stars become available.  This talk will also show that this method can be used with measurements of other macroscopic observables, like the masses and tidal deformabilities, which can (in principle) be measured by gravitational wave observations of binary neutron-star mergers. ",
        "introduction": "Models of stars in general relativity theory are constructed by solving the relativistic stellar structure equations that follow from Einstein's equation.  The most widely used form of these equations, for non-rotating stellar models, was first derived by Oppenheimer and Volkoff~\\cite{Oppenheimer1939}: \\begin{eqnarray} \\frac{dm}{dr} &=& 4\\pi r^{\\,2} \\epsilon, \\label{e:OVmass}\\\\ \\frac{dp}{dr} &=& -(\\epsilon+p)\\frac{m + 4\\pi r^{\\,3} p}{r(r-2m)}. \\label{e:OVpressure} \\end{eqnarray} These equations determine $m(r)$ and $p(r)$, the mass of the star contained within a sphere of radius $r$ and the pressure respectively. These equations are incomplete however, and must be supplemented by giving an equation of state that characterizes the material contained within the star.  In particular, the relationship between the total energy density $\\epsilon$ and the pressure $p$ must be specified.  The addition of the relationship $\\epsilon=\\epsilon(p)$ completes the structure equations, (\\ref{e:OVmass}) and (\\ref{e:OVpressure}), making them a first-order set of ordinary differential equations that can be solved using standard techniques.  These equations are usually integrated starting at the center of the star, where $r=0$, using the boundary conditions $m(0)=0$, to ensure that the stellar model is non-singular, and $p(0)=p_c$ to select the central value of the pressure.  The value of the central pressure $p_c$ can be chosen arbitrarily, with larger $p_c$ generally corresponding to stellar models having larger masses.  The stellar structure equations, (\\ref{e:OVmass}) and (\\ref{e:OVpressure}), are then integrated toward larger values of $r$ until the surface of the star at $r=R$ is reached, where $p(R)=0$.  The total mass of the star is given by $M=m(R)$.  The central pressure $p_c$ is a freely specifiable parameter, so the stellar structure equations determine a one-parameter family of solutions, $m(r,p_c)$ and $p(r,p_c)$, for any given equation of state. In general, the macroscopic observables associated with these stellar models, like the total mass $M$ and radius $R$, also depend on this parameter $p_c$.  Therefore the macroscopic observables associated with any particular equation of state can be thought of as a curve, parameterized by $p_c$, in the space of those observables.  This talk will focus on the particular case of masses $M(p_c)$ and radii $R(p_c)$.  The standard stellar structure problem can be thought of as an abstract map that associates an equation of state (which is in effect a curve though the space of energy densities and pressures) with a curve through the space of observables.  Figure~\\ref{f:Fig1} illustrates what this abstract map looks like for the particular case of the mass-radius curve of observables. \\begin{figure}[!t] \\includegraphics[width=4.4in]{lindblom_Fig1.eps} \\caption{Standard relativistic stellar structure problem can be thought of as an abstract map between an equation of state (a curve in the space of densities and pressures) and a curve of macroscopic observables, like the mass-radius curve. \\label{f:Fig1}} \\end{figure} ",
        "conclusions": "This talk has discussed various aspects of the relativistic inverse stellar structure problem, i.e.  the problem of determining the equation of state of the stellar matter from a knowledge of a curve through the space of macroscopic observables of those stars.  Some recent work on the use of spectral methods was described.  Spectral methods are a very efficient way to represent functions, and the particular spectral expansions described here were shown to be capable of describing a wide variety of theoretical neutron-star equations of state with accuracies of a few percent using spectral approximations with only two or three spectral parameters.  It was also shown that these spectral parameters could be determined by fitting models of the masses and radii or the masses and tidal deformabilities of neutron stars with accuracies that are only somewhat less accurate that the optimal spectral fits to those equations of state.  The spectral solutions to the inverse stellar structure problem described here were based on the use of very high quality measurements of the macroscopic observables that are uniformly distributed across the astrophysically relevant range of neutron-star masses.  It remains to be seen how well these spectral methods perform when more realistic data us used, i.e. data with finite measurement errors and data that is maldistributed in some way.  I plan to investigate those issues at some point in the near future.   \\begin{theacknowledgments} I thank the organizers of the Fifth Leopoldo Garc\\'ia-Col\\'in Mexican Meeting on Mathematical and Experimental Physics for arranging a scientifically interesting and extremely enjoyable conference.  I am grateful for having had the opportunity to participate, to renew old scientific friendships, and to make a number of new ones.  I also thank the Mathematical Sciences Center at Tsinghua University in Beijing, China for their hospitality during the time this manuscript was being written.  The research reported in this talk was supported by a grant from the Sherman Fairchild Foundation and by NSF grants PHY1005655 and DMS1065438. \\end{theacknowledgments} \\vfill\\break"
    },
    "1402/1402.0345_arXiv.txt": {
        "abstract": "We present nuclear spectral energy distributions (SEDs) from 1 to 18 \\micron~of a small sample of nearby, nearly face-on and undisturbed Seyfert galaxies without prominent nuclear dust lanes. These nuclear SEDs probe the central $\\sim$35 pc of the galaxies, on average, and include photometric and spectroscopic infrared (IR) data. We use these SEDs, the clumpy torus models of Nenkova et al. and a Bayesian approach to study the sensitivity of different IR wavelengths to the torus parameters. We find that high angular resolution 8--13 \\micron~spectroscopy alone reliably constrains the number of clumps and their optical depth (N$_0$ and $\\tau_V$). On the other hand, we need a combination of mid- and near-IR subarcsecond resolution photometry to constrain torus width and inclination, as well as the radial distribution of the clouds ($\\sigma$, $i$ and $q$). For flat radial profiles ($q$=0, 1), it is possible to constrain the extent of the mid-IR-emitting dust within the torus ($Y$) when N-band spectroscopy is available, in addition to near-IR photometry. Finally, by fitting different combinations of average and individual Seyfert 1 and Seyfert 2 data, we find that, in general, for undisturbed, nearly face-on Seyferts without prominent nuclear dust lanes, the minimum combination of data necessary to reliably constrain all the torus parameters is J+K+M-band photometry + N-band spectroscopy. ",
        "introduction": "\\label{intro}  Recent success in explaining several properties of the nuclear infrared (IR) spectral energy distributions (SEDs) of Seyfert galaxies has been gathered with the assumption of a clumpy distribution of dust surrounding active galactic nuclei (AGN). Clumpy torus models \\citep{Nenkova02,Nenkova08a,Nenkova08b,Honig06,Honig10b,Schartmann08,Schartmann09,Stalevski12} propose that the dust is distributed in clumps, instead of homogeneously filling the torus volume. The latter authors explain in detail how different model parameters control the shape of the synthetic SEDs, and in some cases compare with observations.  From the observers point of view, our aim is to reproduce nuclear SEDs of nearby AGN with torus models. In turn the goal is to look for trends in the torus model parameters with different observational AGN properties such as luminosity, AGN class and presence or not of broad lines (\\citealt{Mason06,Mason09,Mason13,Mor09,Horst08,Horst09,Nikutta09}, Ramos Almeida et al. 2009, 2011a,b; \\citealt{Honig10b}; Alonso-Herrero et al.~2011, 2012, 2013; \\citealt{Lira13}).   In a series of papers, we performed detailed fits to the nuclear near- and mid-IR (NIR and MIR) emission of Seyfert galaxies using the clumpy torus models of \\citet{Nenkova08a,Nenkova08b} and a Bayesian approach (see below). In the first two works of the series (Ramos Almeida et al. 2009, 2011a; hereafter \\citealt{Ramos09} and \\citealt{Ramos11} respectively), we only fitted subarcsecond resolution photometry from 1--20 \\micron. The main results of these studies are: 1) the NIR/MIR emission of Seyfert galaxies can be reproduced with small torus sizes, ranging from 1 to 6 pc. 2) Seyfert 2 (Sy2) tori generally have larger covering factors and smaller escape probabilities than those of Seyfert 1 (Sy1) galaxies. 3) The classification of a Seyfert galaxy as a Type-1 or Type-2 might depend, not only on torus inclination, but also on the intrinsic properties or that torus.  In the third paper of the series (Alonso-Herrero et al.~2011; hereafter \\citealt{Alonso11}), we combined ground-based subarcsecond MIR spectroscopy and NIR/MIR photometry of 13 Seyfert galaxies at a median distance of 31 Mpc. In general, the photometry+spectroscopy fits resulted in more constrained torus widths ($\\sigma$), torus radial extents ($Y$) and viewing angles ($i$). From the study presented in \\citealt{Alonso11}, we learnt that clumpy torus models provide good fits only for those Seyferts with low--to--moderate amounts of foreground extinction (A$_V\\la$ 5--10 mag). The 8--13 \\micron~spectra of galaxies with very deep silicate features, normally hosted in highly inclined or interacting galaxies and/or showing prominent dust lanes, cannot be reliably reproduced with clumpy models \\citep{Levenson07}. The latter result was confirmed in the recent analysis of ground-based MIR spectra of nearby Seyferts presented in \\citet{Gonzalez13}. The host galaxies of the Seyfert nuclei with the largest nuclear 9.7 \\micron~apparent optical depths show nuclear dust lanes at optical wavelengths, are highly inclined and/or are part of a merger. The dust associated to these galaxy properties is likely contributing to the deep silicate features observed. A similar conclusion was reached by \\citet{Goulding12} for a sample of Compton-thick Seyfert galaxies using data from the Spitzer Infrared Spectrograph (IRS) and probing kpc scales.    The previous works are useful for testing the models and constraining the torus parameters of AGN, considering that with current instrumentation we cannot obtain direct observations of the torus itself. The problem is that, when fitting SEDs with torus models, we normally use the IR data available in the literature for our targets. This implies that, for different objects of a given sample, the data were obtained with distinct filters and instruments, and the sampling of the SEDs may not be the same. Distinct SEDs of the same target can provide different fitting results, and for that reason, it is important to analise how IR observations at various wavelengths constrain the torus model parameters. This analysis is important because, to isolate as much as possible the torus emission, subarcsecond resolution data from ground-based 8-10 m class telescopes or the {\\it Hubble Space Telescope (HST)} are necessary, and their over-subscription factors are very high.  {\\it BayesClumpy} \\citep{Asensio09} is a computer program that can be used for the fast synthesis of SEDs emerging from the clumpy dusty torus models of Nenkova. These fast synthesis capabilities are used in a Bayesian scheme for carrying out inference over the model parameters for observed SEDs. In the latest version of {\\it BayesClumpy}\\footnote{https://github.com/aasensio/bayesclumpy}, the inference can be done either using neural network interpolation (as described in \\citealt{Asensio09}) or multilinear interpolation in the full database. After running different tests, we recommend the latter interpolation to the users for fitting the clumpy torus models of Nenkova. In addition, we have recently incorporated Bayesian adaptive exploration (BAE) in {\\it BayesClumpy} to predict which photometric filter is needed to maximize the expected utility \\citep{Asensio13}. In other words, we used the information already present in the data and clumpy torus models to evaluate the optimum next observation that maximizes the constraining power of the new observed photometric point.  Here we use BAE to study the SED sensitivity (including photometric and spectroscopic IR data) to the six clumpy torus parameters of Nenkova. We also propose the optimum filterset required to constrain these parameters, using average and individual Sy1 and Sy2 SEDs of a small sample of nearly face-on, undisturbed Seyfert galaxies without prominent nuclear dust lanes. Although the Nenkova models are generic and widely used by the community, we note that some of the parameters that describe them might vary from model to model \\citep{Honig06,Schartmann08,Stalevski12}. Besides, distinct models often treat differently the radiative transfer problem. In fact, a given SED can be reproduced with a different set of two distinct model parameters. Thus, the results from SED fitting with torus models should be interpreted in the context of the chosen model only, as the resulting parameters are just a possible description of the geometry and properties of the torus, provided that the data are dominated by torus emission. See \\citet{Honig13} for further discussion on SED fitting caveats. In summary, the results presented here are only valid, in principle, for the clumpy torus models of Nenkova. Throughout this paper we assume a cosmology with H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.27, and $\\Omega_{\\Lambda}$ =0.73. ",
        "conclusions": ""
    },
    "1402/1402.4170_arXiv.txt": {
        "abstract": "Hubble Space Telescope imaging of the galaxy cluster Abell 1689 has revealed an exceptional number of strongly lensed multiply-imaged galaxies, including high-redshift candidates.  Previous studies have used this data to obtain the most detailed dark matter reconstructions of any galaxy cluster to date, resolving substructures ~25 kpc across. We examine Abell 1689 (hereafter, A1689) non-parametrically, combining strongly lensed images and weak distortions from wider field Subaru imaging, and we incorporate member galaxies to improve the lens solution. Strongly lensed galaxies are often locally affected by member galaxies, however, these perturbations cannot be recovered in grid based reconstructions because the lensing information is too sparse to resolve member galaxies. By adding luminosity-scaled member galaxy deflections to our smooth grid we can derive meaningful solutions with sufficient accuracy to permit the identification of our own strongly lensed images, so our model becomes self consistent. We identify 11 new multiply lensed system candidates and clarify previously ambiguous cases, in the deepest optical and NIR data to date from Hubble and Subaru. Our improved spatial resolution brings up new features not seen when the weak and strong lensing effects are used separately, including clumps and filamentary dark matter around the main halo. Our treatment means we can obtain an objective mass ratio between the cluster and galaxy components, for examining the extent of tidal stripping of the luminous member galaxies. We find a typical mass-to-light ratios of $M/L_B = 21 \\pm 14$ inside the $r<1$ arcminute region that drops to $M/L_B = 17 \\pm 8$ inside the $r<40$ arcsecond region. Our model independence means we can objectively evaluate the competitiveness of stacking cluster lenses for defining the geometric lensing-distance-redshift relation in a model independent way. ",
        "introduction": "\\begin{figure*} \\includegraphics[width=8cm]{figs/MassShearArcs_A1689_SLWL_10arcmin.ps} \\includegraphics[width=8cm]{figs/A1689_Gal_Arcs_3p33arcmin.ps}  \\caption{Data set used for the reconstruction. The left panel shows the entire field of view of $10\\arcmin$ where the strong lensing arcs are represented in white and the reduced shear measurements are represented in black. The position of the galaxies and their shapes are represented in light grey. The right panel shows the central $3.33\\arcmin$ region with the arcs and galaxies in non-linear scale to better show the extent of the galaxy haloes in our model.} \\label{fig_data} \\end{figure*}  A fuller exploration of non-parametric cluster lensing is increasingly motivated by new dedicated deep Hubble imaging surveys, with the aim of examining dark matter structures in the least biased way. Multiple sets of lensed images are now typically identified in deep, high resolution images of any cosmologically distant cluster imaged with Hubble, allowing systematic exploration of the cluster dark matter and discovery of the most distant galaxies \\citep{Broadhurst05a,Clowe2006,Coe2010,Coe2012,Coe2013,Zitrin2009,Zitrin2010,Zitrin2011, Zheng2012}.  In practice, secure identification of multiple images need the guidance of a reasonably accurate lens model as even the counter images of large arcs are typically hard to find given the complexities in the central mass distribution of clusters, so that images for a given source are far from symmetrically located.  Furthermore, the uncertain redshifts of faint lensed images means that even when a reliable mass model can be built, counter images are predicted to fall on long, largely radial loci or may fail to be generated at all if the unknown source distance is insufficient. This means that there are often several contending counter images unless internal colors and morphology are sufficiently distinctive. Photometric redshifts, if unambiguous, are very helpful in limiting the selection of counter images that are too faint for spectroscopy. The 16 overlapping broad bands of the CLASH program covering the UV to the NIR, maximises the photometric redshift accuracy possible with Hubble and have provided reliable examples of the most distant galaxies known, as in the case of the $z\\simeq 11$ candidate lensed by MACS0647 \\citep{Coe2013}, where multiple images are identified both photometrically and geometrically.  \\begin{figure*} \\includegraphics[width=5.8cm]{figs/Critical_curve_SL_8000IT.ps} \\includegraphics[width=5.8cm]{figs/Critical_curve_WL_100IT.ps} \\includegraphics[width=5.8cm]{figs/Critical_curve_SLWL_8000IT.ps} \\caption{Solutions (and associated critical curves for a source at redshift z=2) for the cases where only the SL data is used in the reconstruction (left), only the WL data is used in the reconstruction (center) and both, the SL and WL data are used in the reconstruction (right). In all cases only the central 6.66x6.66 arcmin$^2$ region is shown. } \\label{fig_SL_WL_SLWL_CritCurves} \\end{figure*}  To date, the galaxy cluster A1689 remains the best studied cosmic lens with hundreds of magnified images in the central region visible in deep Hubble images \\citep{Broadhurst05a}. Over a hundred of these arcs have been matched to their corresponding background galaxies by several authors and their redshifts estimated \\citep{Broadhurst05a,Halkola2006,Limousin2007,Coe2010}, including several of the brightest highest redshift galaxies known, extending to $z\\simeq 7.6$ \\citep{Frye2002, Bradley2008}. The relaxed appearance of this cluster and the relatively undisturbed optical and X-ray morphology has made this a preferred target for constraining the equilibrium mass profile by several independent means \\citep{Lemze2008,Sereno2013}.  Several studies have used these arcs to reconstruct the mass distribution using the strong lensing data alone \\citep{Broadhurst05a,Diego05b,Halkola2006,Jullo2009,Coe2010} and in combination with weak lensing (or WL hereafter) measurements \\citep{Broadhurst05b,Limousin2007,Umetsu2008} including the use of background red galaxies whose surface density is depleted by lens magnification and independent observationally from weak lensing shear. A1689 has been studied also using higher order derivatives of the lensing potential, like the flexion (see \\citep{Leonard2011} for a recent analysis) The mass profile of A1689 was shown to be very well fitted by the standard NFW profile describing the equilibrium mass distribution expected for collisionless, cold dark matter (CDM) \\citep{Broadhurst05a,Broadhurst05b} but with a concentration that is surprisingly high. Triaxiality of the mass distribution has been explored as a means of boosting projected concentrations, and certainly may be expected to be partially responsible \\citep{Oguri2005, Broadhurst2008,Sereno2013}.     A1689 has been subsequently followed at other wavelengths, allowing lensing to be combined with SZ and X-ray data \\citep{Sereno2013} and also with the dynamics of member galaxy motions via the Jeans equation and via velocity caustics \\citep{Lemze2009}. Multiwavelength science opens the door to new exciting studies since it is no longer sufficient to model the mass or gas separately but instead both have to be integrated in the same model in order to explain the observations. Previous work on A1689 combines HST and {\\it Chandra} and reveals some tension between hydrostatic+lensing reconstruction and other observations  \\citep{Zekser2006,Leonard2007,Umetsu2008,Cain2011, Sereno2011,Lemze2008}. \\cite{Peng2009} finds a discrepant hydrostatic mass based on X-ray data from {\\it Chandra} but \\cite{Riemer2009} find that excluding substructure alleviates or even eliminates the discrepancy. Lensing data combined with X-ray and SZ data have the potential to reveal information not only about the {\\it invisible} dark matter distribution but also about the physical phenomena taking place in the cluster that have to bring the gas pressure and dark-matter-driven gravity to a quasi-equilibrium state. The new Frontier Fields program\\footnote{http://www.stsci.edu/hst/campaigns/frontier-fields/} is now underway to provide the deepest Hubble data ever recorded for massive lensing clusters, further motivating our assumption-free modeling.    Despite exhausting lensing studies of A1689, many arcs still remain unmatched for this cluster. Certainly, many missing counterimages of highly magnified images are too faint to be useful or remain undetected. Others are of too low contrast to be detected within the light of luminous cluster members. Inaccuracy of mass models used to reconstruct the mass distribution is another issue, given the significant variation between published solutions. In \\cite{Ponente2011}, the authors show how very erroneous mass distributions can still reproduce lensing data to high accuracy.  In other words, being able to reproduce the observed arcs is no guarantee that the reconstructed mass distribution is right. The most detailed analysis made using A1689 strong lensing data, and that have been able to match tens of arcs, have relied either on the distribution of member galaxies to guide the models or on parametric models with the inclusion of many parameters depending on the number of substructures adopted.  However, the persistent resistance of some obvious bright lensed images to be matched with other images (when the same models predict another bright counterpart for those lensed images) suggests that these models are still missing fundamental pieces that allows to solve the puzzle. One of the limitations of strong lensing data is that it quickly becomes insensitive to the mass distribution beyond the Einstein radius, especially if the distribution of matter around the center is more or less spherical. Clumps with a significant amount of matter that lie just beyond the Einstein radius might go unnoticed with these parametric models as the model contains enough parameters within the Einstein radius to easily fit the data. Attempts have been made to constrain the matter distribution beyond the Einstein radius of A1689 by combining the weak lensing data with strong lensing data but in all cases (to the best of our knowledge) these {\\it joint} analysis have been made a posteriori where either the density profiles are combined to extract a single density profile \\citep{Umetsu2008} or the SL solution is tested against the WL data for consistency \\citep{Limousin2007}.  In this paper we revisit this cluster to obtain a truly joint solution combining in a single inversion (i.e. not {\\it a posteriori}) the SL and WL data. By doing this, our 2-dimensional model of the mass distribution has to account simultaneously for the multiple lensed systems observed in A1689 and for the shear measurements that extend well beyond the Einstein radius.  \\begin{figure} \\includegraphics[width=8cm]{figs/MagnificationMap_SLWL_8000IT.ps} \\caption{Magnification map in the entire field of view of 10x10 arcmin$^2$ for the SL+WL case. The colors are in log-scale to increase contrast. The two small black regions in the centre correspond to magnifications less than 1 (that have been set to 0 in the log-scale for contrast purposes). The drop in the outer region (buffer zone) is a systematic effect due to the larger cell size in the grid. The features close to the phase transition region between the two grid resolutions are not always to be trusted.} \\label{fig_magnification} \\end{figure} ",
        "conclusions": "We have presented a robust estimation of the dark matter distribution in the cluster A1689. We explore the range of variability of the solutions and identify a region of minimum variance where the solution is stable against changes in the configuration of the data set, the number of iterations, the grid resolution and, assumptions made of the fiducial deflection field. We also identify regions where the results should be taken with more caution.  Our solution can be used to identify additional strongly lensed systems. We identify 11 new systems (candidates) and confirm some of the previous identifications like including the contentious system number 12 where our solution correctly predicts the fourth arclet identified in \\cite{Limousin2007}.  Even though the WL measurements have a typical sampling scale of $1\\arcmin$, through the combination of the SL and WL data sets in the same minimization algorithm we manage to improve upon this resolution beyond the Einstein radius and be sensitive to smaller scales. This allows us to resolve details unseen before in the dark matter distribution around and beyond the Einstein radius, some of which have no obvious correlation with the luminous matter.  At larger radii ($r>2$ arcmin) the sensitivity to smaller details weakens as the SL data set loses its capability to constrain the matter distribution at these distances.  A consequence of knowing the mass distribution of a lens is that it makes it possible to make predictions that can be used in other observations. For instance, we can derive redshifts for sources with unknown redshift by projecting the system back at different redshifts and finding the redshift at which the system {\\it come into focus}. Figure  \\ref{fig_GravLens89} is an example for two of the photometric redshift systems in our sample. For these sources, the lens predicts lower redshifts than the photo-z although the predicted redshift is still consistent with the photo-z.  The mass model can be also used to impose constraints on the cosmological model, for instance through the $f_k$ function as shown in section \\ref{sect_cosmo} although we also show that in order to get competitive constraints one should probably rely on stacking results from multiple clusters.  Finally, our free-form model allows us to determine the mass-to-light ratio of the main galaxies in the cluster. We find ratios that are generally consistent with earlier results found in the literature. However, with the exception of the central galaxy, the galaxies in our model assume the same $M/L_B$ ratio limiting somehow the power of our study. A more detailed study where member galaxies are allowed to take on individual $M/L_B$ ratios will be considered in a future work.   \\begin{figure} \\centerline{ \\includegraphics[width=8cm]{figs/GravLens_Z_system8dotted_system9dashed.eps}} \\caption{Gravitational lensing redshift prediction for systems 8 (dotted) and 9 (dashed) from our reference model. } \\label{fig_GravLens89} \\end{figure}"
    },
    "1402/1402.3750_arXiv.txt": {
        "abstract": "We report on our searches for debris disks around seven relatively nearby radio pulsars, which are isolated sources and were carefully selected as the targets on the basis of our deep $K_s$-band imaging survey. The $K_s$ images obtained with the 6.5\\,m Baade Magellan Telescope at Las Campanas Observatory are analyzed together with the \\textit{Spitzer}/IRAC images at 4.5 and 8.0~$\\mu$m and the \\textit{WISE} images at 3.4, 4.6, 12 and 22~$\\mu$m. No infrared (IR) counterparts of these pulsars are found, with flux upper limits of $\\sim \\mu$Jy at near-infrared ($\\lambda<10\\ \\mu$m) and $\\sim$10--1000\\,$\\mu$Jy at mid-infrared wavelengths ($\\lambda>10\\ \\mu$m). The results of this search are discussed in terms of the efficiency of converting the pulsar spin-down energy to thermal energy and X-ray heating of debris disks, with comparison made to the two magnetars 4U~0142+61 and 1E~2259+586 which are suggested to harbor a debris disk. ",
        "introduction": "Disks are ubiquitous in astrophysical systems at all scales, found almost anywhere an axis of symmetry exists, and play an important role in the appearance and evolution of these systems. For example, they are essential for the creation of planetary systems, formation of young stars, mass-transfer in binary systems, and fueling of the central engines in active galactic nuclei, even though the physical properties of both the disks and their host systems differ dramatically.  Disks may also exist around isolated neutron stars, possibly formed from fallback material after supernova explosions \\citep{lwb91,chn00}. When a massive star dies in a supernova explosion, some amount of the ejecta ($\\leq 0.1 M_{\\odot}$) may ``fallback'' onto the newly formed neutron star \\citep{che89,ww95}. If much of the material is captured by the gravitational field of the neutron star and has sufficiently high angular momentum to prevent direct infall onto this compact star, a disk forms due to interactions between streams of the captured material.  There are a number of observational puzzles relating to radio pulsars that could be explained by disks \\citep{md81,cs08}.  The evolution of pulse periods with time for most pulsars does not actually follow the detailed predictions of magnetic dipole radiation (e.g., \\citealt{liv+07,esp+11}), which could be due to disks \\citep{mph01,cal+13} exerting excess torque on the pulsars' magnetospheres. In addition, some unique phenomena seen in individual pulsars, such as jets and neutron star precession, may arise from disks \\citep{bp04,qia+03}. The first extra-solar planet system was discovered around the radio pulsar B1257+12 \\citep{wf92}. Even though the source is a millisecond pulsar with a complicated evolutionary path, it is generally believed that the planetary system had formed from a debris disk (e.g., \\citealt{ph93,mh01}). There is a type of radio emitting neutron stars, called rotating radio transients (RRATs; \\citealt{mcl+06}), which are characterized by short, periodic radio bursts (2--30 ms) spaced by long intervals (ranging from minutes to hours). The bursts, along with other variations in radio pulses---such as pulse nulling, drifting subpulses, and mode changes, might be caused by dusty material from a surrounding debris disk falling into the magnetosphere of a pulsar \\citep{cs08,sha+13}.  Motivated by the discovery of the planetary system around PSR B1257+12, searches for debris disks around neutron stars have concentrated on old radio pulsars, but with no success (e.g. \\citealt{ff96,koc+02,lf04}). These searches were however limited by the choice of targets (older, less energetic pulsars) and the instrumentation (poor sensitivity and angular resolution). Multi-wavelength searches for fallback disks around a number of young neutron stars have been conducted (e.g., \\citealt{wck06,wkc07}), although these are considered as high-energy objects, namely, central compact objects (CCOs; \\citealt{pst04,gha13}) and magnetars \\citep{td96}, rather than ordinary/traditional radio pulsars. The former are a class of radio-quiet X-ray point sources at the centers of young supernova remnants while the latter are believed to generally have ultra-high surface magnetic fields ($\\sim10^{14}$~G; \\citealt{wt06}).  Using the {\\it Spitzer Space Telescope}, the infrared (IR) counterparts to the magnetars 4U 0142+61 and 1E~2259+586 have been found \\citep{wck06,kap+09}. For the former source, no single spectral model, such as a blackbody or a power law, can fit its combined optical and IR spectral energy distribution (SED). Instead, it is very likely that the IR emission arises from a residual dust disk.  This disk, with estimated mass of $<$10$^{-3} M_{\\odot}$, is heated by strong X-rays from the central pulsar ($L_{\\rm X}\\simeq 10^{36}$ erg s$^{-1}$) and appears bright at IR wavelengths (for example, $L_{\\rm 4.5\\mu m}\\simeq 5\\times 10^{31}$ ergs s$^{-1}$).  The disk's lifetime was estimated to be $>10^6$~yr, assuming that the magnetar has been spun down by mass accretion from the putative disk \\citep{wck06}.  The lifetime significantly exceeds the pulsar age ($\\approx 10^5$ yr), suggesting a supernova fallback origin \\citep{wck06}. While there were only $K_s$-band and \\spz\\ 4.5 $\\mu$m detections of 1E~2259+586, the overall IR emission is similar to that of 4U 0142+61, suggesting the possible existence of a debris disk around this magnetar \\citep{kap+09}.   In order to explore the general existence of similar residual disks around isolated neutron stars, we have carried out disk searches around a few relatively young radio pulsars.  The detectability of the disk around 4U~0142+61 is likely the result of a combination of the very strong X-ray heating of the disk and the relative youth of the object. For radio pulsars, their winds that carry much of their spin-down luminosity and contain energetic particles may heat a debris disk, producing IR emission \\citep{jon07,jon08}. Some fraction of the spin-down luminosity is conceivably assumed to illuminate surrounding disks.  In this paper we report the results of disk searches around seven relatively young pulsars from our ground-based and \\spz\\ observations. The pulsar targets were carefully selected partly on the basis of an initial near-IR $K_s$-band survey. Results from the all-sky survey conducted by the \\textit{Wide-field Infrared Survey Explorer} (\\textit{WISE}) were included. In addition, in order to understand the role of X-ray emission from neutron stars in disk heating (or disk detectability), X-ray flux values or upper limits of the seven pulsars, either from the literature or archival data analysis conducted by us, were also reported.  Below we first describe our target selection in Section~\\ref{subsec:ts}. The observations and data analysis are described in Section~\\ref{sec:obs}, and the results and discussion are presented in Sections~\\ref{sec:res} and \\ref{sec:dis}, respectively.  \\subsection{Target Selection} \\label{subsec:ts} We selected sources from the Pulsar Catalog\\footnote{http://www.atnf.csiro.au/people/pulsar/psrcat/} \\citep{man+05} based on the spin-down luminosity $L_{\\rm sd}$ scaled by distance $d$, $L_{\\rm sd}/d^2$, and picked the top 40 young, energetic pulsars that had not been well observed in the optical/IR bands.  As one might expect near-IR emission from these systems if debris disks are present, we undertook observations in the $K_s$ band using the 6.5-m Magellan telescopes (see Section~\\ref{sec:obs}). However, in no case did we find a near-IR counterpart.  Based on the ground-based observations, we then selected seven pulsars (see Table~\\ref{tab:tgt}) as our targets for \\spz\\ observations. The reasons for the selection are the following.  First, they are relatively nearby, with a typical distance of 4 kpc.  A debris disk with temperature of approximately 1000~K \\citep{wck06} would be bright in the mid-IR at that distance. Second, these targets are not surrounded by synchrotron nebulae that can envelope some of the most energetic pulsars \\citep{krh06,gs06} and could hide mid-IR emission from the pulsar itself. Third and most important, since most pulsars are located primarily in the Galactic plane, the confusion with field stars is a serious problem.  These targets were selected because they are in less crowded fields.  Using the Galactic disk dust map \\citep{sfd98}, we estimated the extinction along each source's direction. The maximum extinction for our targets is approximately $A_V=9$ (for three high Galactic latitude sources, they only have $A_V\\approx 2$). This limits the extinction effect that could alter the apparent stellar colors and cause us to mistake it for a disk.  To our target selection we added one other source: RRAT J1317$-$5759. This object, which just emits sporadic bursts of radio waves, would not be detected simply based on our spin-down luminosity argument. However, the mechanism for the radio bursting is not clearly understood, and may be related to a surrounding debris disk \\citep{cs08}. Given that this source also has a clean field, we selected it as our target for testing this hypothesis. ",
        "conclusions": "\\label{sec:dis}  For the pulsar targets, their $L_{\\rm sd}$ and $d$ (estimated from their DMs and generally have 30\\% uncertainties) are listed in Table~\\ref{tab:tgt}. Except PSR J1317$-$5759, the pulsars have spin-down fluxes of 7--90$\\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$. If a debris disk exists around a pulsar, the energy flux of the pulsar might be intercepted and a fraction of it would be re-radiated by the disk in the IR (e.g., \\citealt{ff96}). Given that the flux range is within an order of magnitude, we use the first pulsar, PSR J0729$-$1448, in Table~\\ref{tab:tgt} as a main example and discuss the implications from our results.  The extinction $A_V$ to each pulsar target was estimated using the three-dimensional extinction maps of the Galaxy \\citep*{dcl03}. The obtained $A_V$ values are given in Table~\\ref{tab:tgt}, which are generally consistent with those initially estimated from the Galactic disk dust map (see Section~\\ref{subsec:ts}). In order to obtain the dereddened flux upper limits, based on which discussions are made below, the reddening laws of \\citet*{sfd98} for the near-IR $K_s$ band data, and \\citet{ind+05} (wavelengths $\\leq$8 $\\mu$m) and \\citet{wd01} (wavelengths $>$8 $\\mu$m) for \\textit{Spitzer} and \\textit{WISE} data were used.  In Figure~\\ref{fig:obs}, the dereddened flux upper limits for PSR J0729$-$1448 are shown. Assuming that a fraction $\\eta$ of the energy flux from this pulsar would be radiated at the IR bands due to the presence of a dust disk [$\\eta(L_{\\rm sd}/4\\pi d^2)=\\nu F_{\\nu}$, where $F_{\\nu}$ is the flux density at a band of frequency $\\nu$], the upper limits have reached $\\eta\\sim 5\\times 10^{-4}$ at 22~$\\mu$m. The $\\eta$ value is nearly as deep as that for the planetary system PSR B1257+12 obtained from \\spz\\ MIPS observations \\citep{bry+06}. However, as pointed out by \\citet{bry+06}, the $\\eta$ limit is still not sufficiently deep since most main-sequence stars with debris disks, which include our solar system, have $\\eta < 10^{-5}$. Even if a dust disk exists, emission would not be strong enough to be detected.  At 4.5 $\\mu$m, our upper limit has reached $\\eta \\simeq 5\\times 10^{-6}$. This ratio is at least one order of magnitude deeper than that of the disk case of the magnetar 4U~0142+61, where $(\\nu_{4.5\\mu}F_{4.5\\mu})/F_{\\rm X}\\sim 10^{-4}$ \\citep{wkc07,wkh07} is used for the comparison because the main energy flux from 4U~0142+61 (as well as 1E~2259+586) is X-ray emission.  In \\citet{wck06}, the IR component in the SED of 4U~0142+61 was modelled to arise from a passive, X-ray irradiated disk (see also \\citealt{vrt+90,phn00}), and as the IR--to--X-ray flux ratio $(\\nu_{4.5\\mu}F_{4.5\\mu})/F_{\\rm X}$ might be typical for disks around X-ray bright neutron stars, the flux ratio has been used for the discussion of disk detectability around different young neutron stars (e.g., \\citealt{wkc07,wkh07}). \\citet{cs08} also have proposed that X-ray emission from a neutron star, due to thermal radiation from its high-temperature surface and/or non-thermal radiation from its magnetosphere, provides heating of a debris disk. To check the possibility that a debris disk would be seen mainly due to X-ray heating, we calculated the $(\\nu_{4.5\\mu}F_{4.5\\mu})/F_{\\rm X}$ flux ratio upper limits for PSRs J07294$-$1448 and B0950+08, which are the only two targets detected with X-ray emission. The ratios are 5.9$\\times 10^{-2}$ and 3.1$\\times 10^{-4}$, respectively, suggesting that our observations were not sufficiently deep. For the pulsar targets without X-ray detection, the X-ray upper limits of B0740$-$28, J0940$-$5428, and J1015$-$5719 are below or nearly equal to 10$^{-3} L_{\\rm sd}$, which is approximately the average fraction of pulsars' spin-down energy emitted at X-rays \\citep{bt97,kar+12}. The flux limits and fluxes suggest that at least half of our pulsar targets (including PSR J0729$-$1448) are unfortunately intrinsically faint at X-ray energies.  In this survey of searching for debris disks around neutron stars, we have focused on relatively young radio pulsars with high spin-down fluxes. While only seven such pulsars were observed with the most sensitive IR telescope due to limited, expensive observing time on a space telescope, the results may suggest that it is difficult to detect any debris disks around radio pulsars, whose existence has been hinted by various pieces of indirect evidence. Considering the detections of possible debris disks around two magnetars, another search focusing on X-ray bright pulsars may be conducted with the future IR telescopes such as the \\textit{James-Webb Space Telescope}. Our survey is also limited to relatively hot debris disks, as the stringent sensitivities were provided by \\spz\\ 4.5 and 8.0~$\\mu$m imaging. In order to complete this survey at multiple wavelengths, searches for cold dust around our pulsar targets should be conducted, given the great capability now available from the Atacama Large Millimeter/submillimeter Array (ALMA) at submillimeter wavelengths. With its sensitivities at least an order of magnitude better than that in the previous searches (e.g., \\citealt{pc94,lww04}), ALMA observations will help draw a full conclusion about our searches for debris disks around isolated radio pulsars."
    },
    "1402/1402.1339_arXiv.txt": {
        "abstract": "A numerical simulation of the Oort cloud is used to explain the observed orbital distributions and numbers of Jupiter-family and Halley-type short-period comets.  Comets are given initial orbits with perihelion distances between 5 and 36 \\AU, and evolve under planetary, stellar and Galactic perturbations for 4.5 Gyr.  This process leads to the formation of an Oort cloud (which we define as the region of semimajor axes $a>1000$ \\AU), and to a flux of cometary bodies from the Oort cloud returning to the planetary region at the present epoch.  The results are consistent with the dynamical characteristics of short-period comets and other observed cometary populations: the near-parabolic flux, Centaurs, and high-eccentricity trans-Neptunian objects.  To achieve this consistency with observations, the model requires that the number of comets versus initial perihelion distance is concentrated towards the outer planetary region.  Moreover, the mean physical lifetime of observable comets in the inner planetary region ($q<2.5$ \\AU) at the present epoch should be an increasing function of the comets' initial perihelion distances.  Virtually all observed Halley-type comets and nearly half of observed Jupiter-family comets come from the Oort cloud, and initially (4.5 Gyr ago) from orbits concentrated near the outer planetary region.  Comets that have been in the Oort cloud also return to the Centaur ($5\\!<\\!q\\!<\\!28$ \\AU, $a<1000$ \\AU) and near-Neptune high-eccentricity regions.  Such objects with perihelia near Neptune are hard to discover, but Centaurs with characteristics predicted by the model (e.g.\\ large semimajor axes, above 60 \\AU, or high inclinations, above $40^\\circ$) are increasingly being found by observers.  The model provides a unified picture for the origin of Jupiter-family and Halley-type comets.  It predicts that the mean physical lifetime of all comets in the region $q<1.5$ \\AU\\ is less than $\\sim$200 revolutions. ",
        "introduction": "\\subsection{Origin of short-period comets}  Explaining the origin of short-period (SP) comets (periods $P<200$ yr) is a long-standing problem. The main difficulty lies in the differences and apparent inconsistency between the respective numbers and orbital distributions of Jupiter-family (JF) and Halley-type (HT) comets. These we classify using the Tisserand parameter $T$ with respect to Jupiter (Carusi et al.\\ 1987), JF comets having $T>2$ (and $P$ usually below 20 years), HT comets having $T<2$ (and $P$ usually between 20 and 200 years). When SP comets are classified this way the number of observed HT comets is found to be less than, or at most comparable to, the number of observed JF comets (see Section \\ref{obs-sp} below).  However, most dynamical theories of their origin from the near-parabolic flux predict a far greater proportion of HT comets (Emel'yanenko and Bailey 1998), with the overall number of observed JF comets conversely being much too large relative to the calculated number (Joss 1973; Delsemme 1973). This discrepancy is associated with the well-known fading problem for long-period comets originating in the Oort cloud and has, at least in part, led to the idea that the two classes of SP comet may have different physical structures and different proximate sources in the present Solar system.  Thus, although there have been many advances in understanding the diverse populations of small bodies in the Solar system, neither a single source dominated by trans-Neptunian objects nor one dominated by the traditional Oort cloud near-parabolic flux at small perihelion distances seems capable of explaining the entire distribution of orbital elements of SP comets. In particular the observed JF comet inclination distribution was recognized to have too many comets at low $i$ relative to the calculated distribution (Duncan et al.\\ 1988; Quinn et al.\\ 1990).  For these reasons, the majority of authors nowadays consider JF and HT comets to be physically as well as dynamically distinct classes, presumably formed in separate regions of the early Solar system and having different dynamical and physical evolutionary histories. Under this viewpoint, JF comets are often regarded as originating largely in the proto-planetary disc beyond Neptune, for example in or close to the Edgeworth-Kuiper belt (EKB). The idea that JF comets might originate in a primordial disc or `belt' of comets located near or beyond the orbit of Neptune was investigated by a number of authors (e.g.\\ Fern\\'andez 1980, 1982; Duncan et al.\\ 1988; Torbett 1989; Torbett and Smoluchowski 1990; Quinn et al.\\ 1990). The discovery of 1992 QB$_1$ (Jewitt and Luu 1993) and of further Edgeworth-Kuiper objects played a pivotal role in theories of the origin of SP comets, and important advances building on this evidence were made in particular by Duncan et al.\\ (1995), Duncan and Levison (1997), Levison and Duncan (1997), and Levison et al.\\ (2001). A key point (Duncan and Levison 1997) was recognition of the potentially important role played by the `scattered' disc, introduced by Torbett (1989) and detected a few years later (Luu et al.\\ 1997), in which it appears that the scattered disc of primordial objects originally formed in the region of the major planets is the principal source of observed JF comets, rather than the EKB. Under the viewpoint of distinct JF and HT classes, HT comets are regarded as objects captured from the Oort cloud (Levison et al.\\ 2001), a structure that would have been produced inevitably as a by-product of planetary, stellar and other perturbations acting on planetesimals or cometary nuclei originally formed by accretion within the planetary region of the proto-planetary disc.  However, a rather unsatisfactory feature of this general picture is the assumption that HT comets coming from the Oort cloud must disintegrate very quickly in order to explain the small number of objects observed (Emel'yanenko and Bailey 1998; Levison et al.\\ 2001). The number of observed inert HT `asteroids' is also very small, and it seems as if the disintegration of a kilometre-size comet nucleus, into presumably an initial trail of much smaller boulder-size objects and then finally dust, must proceed fairly rapidly and lead to eventual extinction of the original comet. On the other hand, dynamical theories appear to require that a high proportion of the JF comet source flux should survive dynamical transfer into the inner Solar system to become active JF comets and that these JF comets should survive for $\\sim$10$^3$ revolutions in the inner Solar system. This difference in the physico-dynamical evolution of the two types of objects is the fading problem for SP comets.  It is probably not unreasonable to assume that comets that formed in different parts of the proto-planetary disc have different physical properties and therefore different lifetimes in the observable region, and it appears that this idea has become very deeply rooted.  What is missing, however, is direct observational evidence to support the idea of two qualitatively distinct types of SP comet, correlating with dynamical class. Thus, present theories of the origin of SP comets rely on a poorly understood fading hypothesis to accommodate the observations, and there is no satisfactory physico-dynamical explanation as to why two very different types of SP comet should exist and yet appear observationally almost indistinguishable.  Indeed, although comets show a very diverse range of properties, covering a very broad range of sizes, densities, dust-to-gas ratios and so on, there is as yet no compelling observational evidence for the expected bimodality of physical characteristics corresponding to HT versus JF dynamical class (Lamy et al.\\ 2004).  In this work, whilst recognizing that comets may have different physical properties depending for example on their sizes or where they might have formed in the proto-planetary disc, we present a model for the common origin and evolution -- from the Oort cloud -- of the majority of comets in the Solar system.   \\subsection{Role of the Oort cloud}  We define the Oort cloud as the region containing objects with semimajor axes $a>10^3$ \\AU\\ (i.e., objects from the Oort cloud have at some point during their evolution reached $a>10^3$ \\AU).  This definition is consistent with those used by other authors; e.g.\\ Wiegert and Tremaine (1999) and Rickman et al.\\ (2008) used similar values of semimajor axis, i.e. $a\\simeq \\mbox{1--3}\\times 10^3$ \\AU, to define the inner boundary of the Oort cloud.  Dones et al.\\ (2004) introduced a further restriction, namely that the maximum value of perihelion distance $q$ during an object's orbital evolution should exceed 45 \\AU\\ for it to be counted as an `Oort cloud' object.  However, objects with $a \\gtsimeq 10^3$ \\AU\\ spend nearly all their time at large heliocentric distances, whatever their value of $q$, and therefore in the Oort cloud.  In this paper we have chosen to define Oort-cloud objects solely according to $a$ because the influence of stellar and Galactic perturbations is determined mainly by $a$ for near-parabolic orbits. It has been shown (e.g.\\ Emel'yanenko 2005) that the dynamical pathways by which objects with $a>10^3$~\\AU\\ reach the planetary region are different from those of typical trans-Neptunian objects (TNOs). While the evolution of TNOs is largely determined only by planetary perturbations, stellar and Galactic perturbations play a more substantial role in the process that drives the perihelia of objects with $a>10^3$ \\AU\\ towards and through the planetary region, regardless of their previous $q$.  In the present paper (Section \\ref{integ} onwards) we numerically integrate a much larger number of objects than in the Oort cloud model of Emel'yanenko et al.\\ (2007), in particular to obtain statistically significant numbers of SP comets captured from the Oort cloud and allow a comparison of the model SP numbers and orbital distributions with the corresponding distributions of observed HT and JF comets.  First (Section \\ref{obs-com}), in order that our model parameters can be constrained by observations, we assess the known characteristics of the various populations of cometary bodies. ",
        "conclusions": "We have developed a model of the origin and evolution of the Oort cloud which is consistent with the basic observed orbital distributions of comets, Centaurs and high-eccentricity trans-Neptunian objects. Rather than requiring intrinsically different fading properties for Jupiter-family and Halley-type short-period comets, the model instead adopts the hypothesis that the physical lifetime of objects as active comets in the inner planetary region at the present epoch is a function of their initial perihelion distance in the early Solar system, and is the same for both JF and HT comets.  The observed JF and HT populations also constrain the initial distribution of objects versus perihelion distance.  Our results show that:  \\begin{enumerate}  \\item The mean physical lifetime of comets is $\\ltsimeq$200 revolutions in the region $q<1.5$ \\AU. This implies a significant cometary contribution to the distribution of small bodies (`boulders' and dust) making up the near-Earth interplanetary complex.  \\item No model in which the initial number of comets is a decreasing function of their initial perihelion distance $q_0$ in the early Solar system can explain the present observed distribution of short-period comets.  \\item Models in which the initial distribution of objects versus perihelion distance is concentrated more towards the outer planetary region, and in which their present active physical lifetime is an increasing function of $q_0$, are consistent with the present orbital distributions and numbers of both HT and JF comets.  \\item Essentially all the observed HT comets and nearly half the observed JF comets come from a proximate Oort-cloud source (i.e.\\ have experienced orbits with $a > 10^3$ \\AU).  The remaining $\\sim$50\\% of observed JF comets come from the {\\em observed\\/} near-Neptune high-eccentricity (NNHE) population, a dynamically unstable region in which the cometary numbers decline by 95\\% over 4 Gyr.  In addition, more than 90\\% of all Centaurs ($5\\!<\\!q\\!<\\!28$ \\AU, $a<1000$ \\AU) come from the Oort cloud.  \\item The model predicts that there is a significant Oort-cloud contribution to the NNHE population.  The number of such objects is comparable to the debiased number of objects already discovered in the NNHE region, but they are still undetected owing to observational biases (e.g.\\ considering large semimajor axes or high inclinations).  \\end{enumerate}"
    },
    "1402/1402.3616_arXiv.txt": {
        "abstract": "We discuss the results of cross-correlating catalogs of bright X-ray binaries with the Yale Southern Proper Motion catalog (version 4.0). Several objects already known to have large proper motions from Hipparcos are recovered.  Two additional objects are found which show substantial proper motions, both of which are unusual in their X-ray properties.  One is IGR J17544-2619, one of the supergiant fast X-ray transients.  Assuming the quoted distances in the literature for this source of about 3 kpc are correct, this system has a peculiar velocity of about 275 km/sec -- greater than the velocity of a Keplerian orbit at its location of the Galaxy, and in line with the expectations formed from suggestions that the supergiant fast X-ray transients should be highly eccentric.  We discuss the possibility that these objects may help explain the existence of short gamma-ray bursts outside the central regions of galaxies. The other is the source 2A~1822-371, which is a member of the small class of objects which are low mass X-ray binaries and long (i.e. $>$ 100 millisecond) X-ray pulsars.  This system also shows both an anomalously high X-ray luminosity and a large orbital period derivative for a system with its orbital period, and some possible indications of an eccentric orbit. A coherent picture can be developed by adding in the proper motion information in which this system formed in the Perseus spiral arm of the Galaxy about 3 Myr ago, and retains a slightly eccentric orbit which leads to enhanced mass transfer. ",
        "introduction": "The mechanisms by which supernovae explode have been a topic of recent debate (e.g. Belczynski et al. 2012).  X-ray binaries represent one of the few classes of systems in the Universe in which the end results of supernova explosions can be explored, and the mass distribution of compact objects has been proposed as a key test of the supernova mechanism.  The systems can be readily identified from their X-ray emission, allowing them to be selected easily for further study.  The masses of the compact objects can be estimated from the radial velocity motions of the donor stars, along with estimates of their inclination angles from eclipses, in an ideal case, or ellipsoidal modulations (see Cantrell et al. 2010 for the state of the art) or polarimetry (Brown et al. 1978; e.g. Dolan \\& Tapia 1989). Astrometric wobble has been seen in Cygnus X-1 (Reid et al. 2011), but not yet at a level where it could be used to make precise measurements of a system inclination angle.  Spaced-based astrometric surveys such as Gaia or the proposed SIM mission do have the potential to measure the inclination angles of wide, nearby X-ray binaries (e.g. Tomsick et al. 2009; Tomsick \\& Muterspaugh 2010).  Along with the mass distribution of compact objects, it is interesting to understand their space velocity distribution.  A few different mechanisms have been proposed for giving X-ray binaries (and, indeed, isolated neutron stars) space velocities far in excess of those of their massive star progenitors.  The loss of mass in supernovae on timescales fast compared with the orbital period means that the mass lost in the supernova will carry away linear momentum relative to the center of mass of the binary system, meaning that the binary itself will receive an equal an opposite momentum kick (see e.g. Blaauw 1961).  These kicks are sometimes referred to as Blaauw kicks, or as symmetric kicks, since they are symmetric about the supernova progenitor's center of mass.  A different class of supernova kicks have also been proposed -- asymmetric kicks, which are, as suggested from the name, asymmetric about the center of mass of the supernova progenitor.  A variety of theoretical mechanisms exist for producing such kicks -- hydrodynamic kicks, due to actual asymmetries in the mass ejection (e.g. Burrows \\& Hayes 1996), interactions between neutrinos and a strong magnetic field (e.g. Arras \\& Lai 1999), and electromagnetic (e.g. Harrison \\& Tademaru 1975) -- see Lai et al. (2000) for a review.  From relatively early on in the study of pulsars, it was found that they had typical space velocities of order 100 km/sec or more (Trimble \\& Rees 1971; Lyne et al. 1982) -- much larger than the space velocities of the OB stars which are their progenitors.  In some cases, these high space velocities might be explained as a ``Blauuw kick''.  A variety of lines of evidence, mostly from observations of X-ray binaries, indicate that some additional kick component is likely to be present (see Podsiadlowski et al. 2005 for a discussion of the array of evidence).  In principle, hypervelocity stars can also be produced in dynamical interactions -- for example with the supermassive black hole in the Galactic Center (e.g. Hills 1988; Brown et al. 2006; Rossi et al. 2013), although the velocities will then always be radial, away from the Galactic Center.  Understanding the space velocities of X-ray binaries gives one route to estimating the kick distribution of neutron stars at birth. Understanding the kick velocities in X-ray binaries is important for understanding their formation and evolution, the retention of neutron stars in globular clusters (e.g. Pfahl et al. 2002; Smits et al. 2006), and for making connections between the populations of high mass X-ray binaries and gravitational wave sources (e.g. Belczynski et al. 2011, 2013).  In this paper, we report on the discovery of large proper motions from two X-ray binaries, the low mass X-ray binary 2A~1822-371, and the supergiant fast X-ray transient IGR~J17544-2619.  In both cases, the high energy properties of the systems are unusual, and large space velocities may help to understand the systems. ",
        "conclusions": "The source IGR~J17544-2619 was first detected with INTEGRAL in a burst of emission that lasted about two hours (Sunyaev et al. 2003), and is a member of the class of objects called supergiant fast X-ray transients (SFXTs).  The SFXTs are a class of objects discovered mostly with INTEGRAL (Sguera et al. 2005) which show extremely bright flares in the hard X-rays on timescales of minutes.  For an excellent review on the topic, we refer the reader to Sidoli (2013), but we briefly summarize some of the key properties of these systems.  About half of the SFXTs show rapid periodicities consistent with the rotation period of a neutron star (Sidoli 2013), and, by definition, all members of the class have supergiant donor stars (e.g. Chaty 2010).  Many of the systems show modulation on the apparent orbital periods of the systems (e.g. Sidoli et al. 2007).  In the Corbet diagram of spin period plotted versus orbital period, the SFXTs share paramter space with both the Be X-ray binaries and the traditional persistent supergiant wind-fed systems (Sidoli 2013).  There remains debate about whether the systems have stellar winds with strong equatorial components (e.g. Sidoli et al. 2007), are clumpy (e.g. in 't zand 2005; Oskinova et al. 2012), or perhaps both.  Gated accretion due to high magnetic fields and slow rotation periods is an additional possibility (e.g. Grebenev \\& Sunyaev 2007; Bozzo et al. 2008).  The full set of system parameters of IGR~J17544-2619 are well described in Drave et al. (2012).  It has been shown to be associated with the infrared counterpart 2MASS J17542527-2619526 (in 't Zand 2005).  Optical and infrared studies have further found that its donor star is an O9Ib star, with a mass of 25-28 $M_\\odot$ (Pellizza et al. 2006) at a distance of about 3.6 kpc (Rahoui et al. 2008).  Its orbital period has been identified by Clark et al. (2009) on the basis of INTEGRAL timing to be 4.926$\\pm$0.001 d, and its pulse period has been found by Drave et al. (2012) to be 71.49$\\pm$0.02 s.  Drave et al. (2012) show that the system lies near the supergiant persistent X-ray binaries in the Corbet diagram, and that the chief difference relative to those systems is likely to be a higher eccentricity for the binaries (see also in 't Zand 2005).  In light of the expectation that these systems might have large eccentricities, measurements of their proper motions is especially interesting, as the same processes that induce eccentricity in previously circular binaries will tend to impart momentum to the binary's center of mass.  These processes can include the loss of a substantial amount of mass during the supernova explosion -- the Blauuw kick -- and an asymmetric kick applied to the compact object during or very shortly after the supernova explosion.  At least one runaway high mass X-ray binary is known, LS 5039 (Ribo et al. 2002).  That system has a short orbital period, and a high measured eccentricity, so it is quite similar in system parameters to IGR~J17544-2619.  On the other hand, it is thought to have a more rapidly rotating neutron star, such that the radio pulsar mechanism works in that system, and its high energy emission is produced by a shock between the stellar wind and the pulsar wind in the system (e.g. Dubus 2006).   A useful starting point in understanding an X-ray binary system is to determine whether its observed kinematical properties could be explained only as the result of a Blauuw kick.  Nelemans et al. (1999) work out the system velocity expected due to symmetric mass loss in a supernovae explosion as: \\begin{equation} v_{sys} = 213 \\left(\\frac{\\Delta{M}}{M_\\odot}\\right) \\left(\\frac{m}{M_\\odot}\\right)\\left(\\frac{P_{re-circ}}{{\\rm day}}\\right)^{-\\frac{1}{3}}\\left(\\frac{m+M_{rem}}{M_\\odot}\\right)^{-\\frac{5}{3}} {\\rm km/sec}, \\label{Nelemans} \\end{equation} in a form which is convenient, but appropriate only for low mass X-ray binaries where there is sufficient time for the system to re-circularize.  They had already derived that $P_{re-circ} = P_{postSN}(1-e^2)^{3/2}$, which is more appropriate to the case of a high mass X-ray binary.  We can then re-write equation \\ref{Nelemans} as: \\begin{equation} v_{sys} = 213 \\left(\\frac{\\Delta{M}}{M_\\odot}\\right) \\left(\\frac{m}{M_\\odot}\\right)\\left(\\frac{P_{postSN}}{{\\rm day}}\\right)^{-\\frac{1}{3}}\\left(1-e^2\\right)^{-\\frac{1}{2}} \\left(\\frac{m+M_{rem}}{M_\\odot}\\right)^{-\\frac{5}{3}} {\\rm km/sec} \\end{equation}  Furthermore, $e = \\frac{\\Delta{M}}{M_{rem}+m}$ -- so it can be seen immediately that for increasing $\\Delta{M}$, both $e$ and $v_{sys}$ increase monotonically.  Furthermore, we can re-write the equation now to remove $e$, getting:  \\begin{equation} v_{sys} = 213 \\left(\\frac{\\Delta{M}}{M_\\odot}\\right) \\left(\\frac{m}{M_\\odot}\\right)\\left(\\frac{P_{postSN}}{{\\rm day}}\\right)^{-\\frac{1}{3}}\\left(1-\\frac{\\Delta{M}}{M_{0,t}}\\right)^{-\\frac{1}{2}} \\left(\\frac{m+M_{rem}}{M_\\odot}\\right)^{-\\frac{5}{3}} {\\rm km/sec} \\label{symmetric} \\end{equation} where $M_{0,t}$ is the initial total mass of the system.  In the case of asymmetric kicks, the situation can be more complicated, since they can counter-act the symmetric kicks (e.g. Kalogera 1998), but generally speaking, large eccentricities and large space velocities are expected to be seen in the same systems.  Starting from equation \\ref{symmetric}, we can estimate the amount of mass loss that must have taken place.  If we take the donor star mass to be 25 $M_\\odot$, and the accretor to be a 1.4$M_\\odot$ neutron star, we find that the systemic velocity will be about $13 \\frac{\\Delta{M}}{M_\\odot}(1-e^2)^{-1/2}$ km/sec.  If the system is assumed to be at a distance of 3.6 kpc, then its tangential velocity relative to its Local Standard of Rest is about 275 km/sec out of the Galactic Plane, {with a range of 160-320 km/sec for the rand of possible distnaces from 2.1-4.2 kpc given in Rahoui et al. (2008). The distance uncertainties dominate over the proper motion uncertainties.}  Its mean systemic radial velocity is about 47 km/sec (Nikolaeva et al. 2013).\\footnote{Nikolaeva et al. (2013) also suggest that the orbital period may be about 12 days for this source, on the basis of 27 spectra.  This analysis gives a minimum mass for the compact object of 2.8 $M_\\odot$, suggesting either that the neutron star is, by far, the heaviest neutron star measured to date, or, more likely, that the radial velocity curve has not yet been sampled well enough, given that fitting an eccentric orbit requires 6 free parameters.  We prefer the latter interpretation, largely on the basis of the very convincing X-ray modulation from the source found by Clark et al. (2009), and the finding that the masses of the neutron stars in high mass X-ray binaries tend to be well-clustered around 1.4 $M_\\odot$ (Thorsett \\& Chakrabarty 1999).  The mean radial velocity should be robust to uncertainties in the orbital solution.}  The system is thus just within the bounds possible for producing a neutron star from a red supergiant without an asymmetric kick being required (see e.g. Woosley et al. 2002).  The eccentricity, at least before any tidal circularization processes have taken place, would be about $13/39.4$, or about 1/3.  At the present time, no concrete measurement exists for the eccentricity of this system.\\footnote{If the Nikolaeva et al. (2013) results are taken at face value then the eccentricity is about 0.44.}  If the eccentricity is small, it may be because the binary has become partly circularized.  Both in 't Zand (2005) and Drave et al. (2012) have suggested that the major difference between SFXTs and the more traditional persistent supergiant X-ray binaries is that the SFXTs have higher orbital eccentricities.  We thus do expect that the binary will not have been completely circularized.  Some models of stellar evolution suggest that neutron stars can be produced from progenitors with quite high initial masses that lose a lot of mass in a Wolf-Rayet phase (Woosley et al. 2002).  If the neutron star were produced from a massive star progenitor, which underwent considerable mass loss in a Wolf-Rayet phase, then it is likely that the progenitor was less than $14 M_\\odot$ before the supernova (Woosley et al. 2002), and that an asymmetric kick was needed to give the system its observed proper motion.  Given the considerable uncertainties, in mass loss rates from massive stars, supernova explosions, and whether this system has mass transfer from the progenitor star of the neutron star to the current donor star, it is difficult to determine without detailed calculations that go beyond the scope of this paper whether the system must have had a natal kick, or whether all its properties can be explained by a simple impulsive mass loss.  The essential point is that the large proper motion is in agreement with the idea that the system at least formed with a large eccentricity.  We can now look at circularization timescales of X-ray binaries to make some estimate of for how long the eccentricity is likely to have remained.  The theory of circularization of binaries is generally tested by observations of open clusters (see e.g. Verbunt \\& Phinney 1995).  However, if, as suggested by Mathieu et al. (1992) that binaries with such short orbital periods ordinarily circularize on the pre-main sequence, then there is no empirical handle on how quickly such systems should circularize.  Most theoretical work has focused on the circularization of binaries with cooler stars (e.g. Zahn 1977; Claret \\& Cunha 1997; Goodman \\& Dickson 1998), so there is not currently even much theoretical work to use to test whether significant circularization should take place in a mass transferring binary with a supergiant donor star.  We can attempt to bound the system's age in a few ways.  First, we can see if we can trace its motion back to any of the star forming regions nearby in the catalog of Avedisova (2002).  The position matches up well with that of region 3.270-0.102 but the distance for that star forming region is given by Downes et al. (1980) to be 18 kpc, so the association is unlikely.  In the absence of a birth site, we can make a crude estimate of the age, on the basis that the system is moving away from the Galactic Plane at about 10 msec/year, and is unlikely to have formed more than about 50 pc from the Galactic Plane.  This yields an age of a few hundred thousand years, so if substantial circularization has taken place, then we immediately learn something interesting about circularization of supergiant binaries.  On the other hand, if we take its high eccentricity as a given based on its SFXT nature, we can use it as a probe of circularization in binaries. Then we can see that it is moving toward the nearest supernova remnant, and is located several degrees from the next nearest supernova remnant (Green 2009), making an association with a supernova remnant unlikely and thus making the system age likely to be larger than the typical lifetimes of supernova remnants.  It is thus unlikely to be much younger than about $10^4$ years.  Further searches for associations (e.g. from the Vista Variables in the Via Lactea project -- Minniti et al. 2010) with young star associations may potential help locate the birth site of this X-ray binary.  \\subsection{Runaway high mass X-ray binaries and the short GRB problem}  The short gamma-ray bursts are thought to be produced by mergers of compact objects, either between two neutron stars or between a neutron star and a black hole (e.g. Paczynski 1986).  As such, high mass X-ray binaries represent excellent candidate progenitors for these events. Furthermore, many of these short gamma-ray bursts have been located not in the centers of galaxies, near the location of maximum star formation, but rather in the outskirts of galaxies (e.g. Gehrels et al. 2005).  The discovery of at least one such object with a peculiar velocity of the same order as the velocity of the Galactic rotation curve helps point the way toward finding the actual progenitors of the subclass of short GRBs taking place outside of the main bodies of galaxies.  In particular, the second velocity kick that will be applied at the time of the second supernova explosion could easily be large enough to unbind the system from the Galaxy without unbinding the binary itself.  The cleanest sample of high mass X-ray binaries whose kinematic properties are well understood is the Hipparcos sample (Chevalier \\& Ilovaisky 1999 -- CI99).  Hipparcos made measurements of 17 HMXBs, including 4 with supergiant donor stars and 13 with Be donor stars (CI 99).  The four supergiant systems are all found to have peculiar tangential velocities of less than 100 km/sec.  The largest quoted tangential velocity is that of Cygnus X-1; the proper motion of Cyg X-1, is, in fact, very similar to that of the Cygnus OB3 association, so that it is unlikely to have formed with a large velocity kick (Mirabel \\& Rodrigues 2003).  The other high mass X-ray binaries studied with Hipparcos have tangential velocities that are larger than the typical velocity dispersion of massive stars, but considerably smaller than the Galactic escape velocity, or even the Galactic rotational velocity.  The Be X-ray binaries measured by Hipparcos have even lower system tangential velocities -- typically about 10 km/sec. That the Be X-ray binaries were formed with relatively small velocity kicks is not surprising given that very large kicks would have unbound most of them.  Developing an understanding of these objects is important, as well, for understanding where gravitational wave sources are likely to be located and determining strategies for conducting electromagnetic follow-up of them.  At the present time, it is not clear what will be the eventual fate of IGR J17544-2619 -- but with continued improvement in understanding its orbital parameters it should be possible to conduct binary evolution simulations (e.g. Belczynski et al. 2006;2008) of it to estimate the probability it will eventually turn into a double neutron star which merges in less than a Hubble time.  \\subsection{Linear momentum} { The linear momentum of the system has a large uncertainty due to the distance uncertainty.  Taking the best estimate distance and the best estimate for the systemic radial velocity, the space velocity of the system is about 280 km/sec.  The donor star mass has been estimated to be 25 $M_\\odot$ (Rahoui et al. 2008), so the systemic mass should be about 26-27 $M_\\odot$.  The total momentum is then equivalent to that of a pulsar with a 5300 km/sec space velocity.  This value is far in excess of the largest space velocities seen from isolated pulsars (Hobbs et al. 2005).  It thus requires that either there is a population of pulsars which has not been detected because pulsars born with such a high velocity quickly escape the Galaxy, or, more likely, that the kick comes primarily from the Blaauw mechanism.  If the kick comes primarily from the Blauuw mechanism, then we can establish that there must have been dramatic mass loss during the supernova that produced the neutron star.  Given that the system must also be quite young, an interesting implication is that the donor star may still be polluted with supernova ejecta.  Abundance studies of the donor stars in SFXTs may thus be an interesting topic for future research.}     We have discussed the matches between the catalogs of bright X-ray binaries and the Yale Southern Proper Motion Survey 4.0 catalog.  We have found two interesting binaries which are both unusual in their X-ray properties, and show large proper motions away from the Galactic Plane, and we have tabulated their three-space velocities in Table \\ref{vel_table}.  In both cases, the X-ray data are suggestive of eccentricities created at the same time the systems received their large peculiar velocities."
    },
    "1402/1402.7086.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small With the increasing number of exoplanets discovered, statistical properties of the population as a whole become unique constraints on planet formation models provided a link between the description of the detailed processes playing a role in this formation and the observed population can be established. Planet population synthesis provides such a link. The approach allows to study how different physical models of individual processes (\\textit{e.g.,} proto-planetary disc structure and evolution, planetesimal formation, gas accretion, migration, etc.) affect the overall properties of the population of emerging planets. By necessity, planet population synthesis relies on simplified descriptions of complex processes. These descriptions can be obtained from more detailed specialised simulations of these processes. The objective of this chapter is twofold: 1) provide an overview of the physics entering in the two main approaches to planet population synthesis and 2) present some of the results achieved as well as illustrate how it can be used to extract constraints on the models and to help interpret observations. \\\\~\\\\~\\\\~}%leave this in to get the correct vertical space after the abstract ",
        "introduction": "} \\bigskip  The number of known exoplanets has increased dramatically in recent years (see, e.g., www.exoplanet.eu; \\citealt{schneider:2011}). At the time of this writing, over 1000 confirmed exoplanets were known, mostly found through precise radial velocity surveys. Additionally, there are more than 3600 transiting candidate planets found by the {\\it Kepler} satellite (see {\\it e.g.,} www.kepler.nasa.gov). All these detections have revealed that planets are quite common and that the diversity of existing systems is much larger than was expected from studies our own Solar System. Finally, with increasing numbers of planets, search for correlations and structures in the properties of planets and planetary systems becomes increasingly meaningful. The correlations and structures have pinpointed the importance of complex interaction processes taking place during the formation stages of the planets (e.g. planetary migration).  These insights were essentially gained by the fact that, for the first time, a  large set of planets was available to study and statistical analysis became possible. The analysis of the characteristics of an ensemble of objects as well as of the differences between individual objects is a standard approach in astrophysics and has been applied successfully in a number of areas (\\textit{e.g.,} galactic evolution).  Planet population synthesis in the context of the core accretion scenario has been pioneered by \\citet{idalin:2004a} in an effort to develop a deterministic model of planetary formation allowing a direct comparison with the observed population of exoplanets. They presented formation models for planets orbiting solar-type stars but neglected the effect of type I migration on the basis that its efficiency was poorly determined. They simulated the mass-semi-major axis distribution of planets for stars of different metallicity and masses \\citep{idalin:2004b, idalin:2005}. \\citet{burkertida:2007} applied this model to discuss a potential period gap in the observed gas giant distribution orbiting stars more massive than the sun. \\citet{currie:2009} discussed the same issue with a similar model. \\citet{paynelodato:2007} discussed planets orbiting brown dwarfs using the model by \\citet{idalin:2004b}.  In \\citet{idalin:2008a}, type I migration was incorporated, using a conventional isothermal formula \\citep[e.g.][]{tanakaetal:2002} with an efficiency factor that uniformly decreases the migration speed, because the predicted formation efficiency of gas giants with the full strength of the migration is too low to be consistent with observations and to allow for uncertainties in the theoretically derived migration speed. \\citet{idalin:2008b} considered the effect of a potential migration trap due to a snow line. With a similar model, \\citet{migueletal:2011a, migueletal:2011b} studied a dependence on initial disc models. \\citet{mordasinietal:2009a} developed a model based on more detailed calculation of gas envelope contraction, disc evolution and planetesimal dynamics. Type I migration was treated in a similar way to \\citet{idalin:2008a}.  They applied their model   \\citep{mordasinietal:2009b} for statistical comparisons with the then known population of giant extrasolar planets.  \\citet{alibertetal:2011} studied the impact of the stellar mass on planetary populations while \\citet{mordasinietal:2012c} investigated how important properties of the proto-planetary disc (mass, metallicity, lifetime) translate into planetary properties. In an attempt to further couple formation models to the major observable physical characteristics of a planet (besides mass and semi-major axis also radius, luminosity and bulk composition) \\citet{mordasinietal:2012} added  to the model the long-term evolution of the planets after formation (cooling and contraction). This enabled statistical comparisons with results of transit and (in future) direct imaging surveys, and in particular  the result of the Kepler satellite \\citep{mordasinietal:2012b}.  Since planet-planet scattering is a chaotic process, including such effects in an otherwise deterministic calculations was not easy. \\citet{alibertetal:2013} incorporated a full N-body integrator with collision detections  in order to simulate planet-planet interactions. While this approach adds a significant computational burden, it has the advantage to handle all dynamical aspects (including resonances) correctly. (Note that \\citet{bromleykenyon:2006, bromleykenyon:2011, thommesetal:2008, hellarynelson:2012} also developed hybrid N-body simulations, although they did not present much statistical discussions of predicted planet distributions). \\citet{idalin:2010} and  \\citet{idaetal:2013} took another approach in which the planet-planet interactions (scatterings, ejections, collisions) are treated in a Monte Carlo fashion calibrated by numerical simulations.  Although relatively complicated multiple steps are needed for the Monte Carlo method to reproduce accurate enough predictions for the statistical purposes of population synthesis, it is much faster than direct N-body simulations.  \\bigskip ",
        "conclusions": "\\bigskip  As demonstrated in this chapter, planet population synthesis is based on a physical description of a large number of processes playing a role in the formation of planets. By necessity, this description is simplified in order to keep the problem tractable and to allow the simulation of the formation of planets for a large set of initial conditions. We have presented the two different approaches that have been most used in the literature. These two sets of prescriptions differ essentially by the degree of simplification being adopted and/or the use of fitting formulas stemming from more complete and detailed calculations of individual processes. While differing in these aspects, both approaches aim at developing a self-consistent model in which the relevant processes operate on their proper timescale. This is necessary to account for the numerous feedbacks occurring between the various processes (see Fig.~\\ref{fig:schematic}).  The detailed physical understanding of the many aspects of planet formation from the early condensation of solids and the formation of planetesimals to the runaway accretion of gas in the late stages of giant planet formation taking place within a time evolving proto-planetary disc is essential. As illustrated in the relevant chapters of this book, some aspects can be studied observationally and some only theoretically essentially by means of large-scale numerical simulations of growing physical accuracy. While all these efforts are essential to our understanding, they are all restricted in either spatial or temporal dimensions.  The goal of population synthesis is to put these space and/or time snapshots together in order to obtain a full picture and a complete history of planet formation from the initial proto-planetary disc to the observed planetary systems. This is essential for at least two reasons: 1) planet formation is not directly observed, only initial conditions (proto-planetary discs) and end-products (planetary systems) are, the link between the two must be provided by theory at least for now; 2) the diversity of characteristics of the ensemble of exoplanets sets statistical constraints on formation models.  In the end, population synthesis is similar to idealised models developed by theorists in an attempt to understand the basic behaviour of complex physical systems. The success of the approach does not lie in having all aspects described in detail but to include the key physical processes and their respective feedback. Hence, population synthesis relies on detailed numerical simulations to provide sufficient understanding to allow the derivation of simplified description that capture the essence of the phenomenon. In return, the approach allows visualising the effect on the ensemble population of planets of a given physical description of individual processes (\\textit{e.g.} disc structure, migration, opacity) which, when compared to observation, allows setting constraints on the models.  We have shown in this chapter that this approach has been useful in the past to identify key problems in the theoretical descriptions on key processes active in planet formation (\\textit{e.g.} planetary migration). We also have pointed out a number of areas where further detailed modelling is needed. Among those, we highlight the following: 1) disc structure and evolution and the associated transport of gas and solids; 2) the formation of planetesimals and the resulting chemical composition and size distribution as a function of distance to the star;  3) migration is still too fast even taking into account the non-isothermal effects associated with the co-rotation torques; 4) the gas flow through the gap is still not clearly established for the accurate determination of gas giant planets' asymptotic masses. The predictive power of population synthesis will rest on the progress that will be achieved in the future in the physical understanding of these processes.  \\bigskip \\noindent\\textit{Acknowledgements.} WB would like to acknowledge partial support from the Swiss National Science Foundation. SI acknowledges support from the JSPS grant. CM acknowledges support from the Max Planck Society through the Reimar-L\\\"ust Fellowship. YA acknowledge the support of the European Research Council under grant 239605. The authors have been supported by the International Space Science Institute, in the framework of an ISSI Team. DL acknowledge support from NASA, NSF, and UC/Lab fee grants.   \\bigskip %\\newpage"
    },
    "1402/1402.5340_arXiv.txt": {
        "abstract": "We developed a {\\sl coronal non-linear force-free field (COR-NLFFF)} forward-fitting code that fits an approximate {\\sl non-linear force-free field (NLFFF)} solution to the observed geometry of automatically traced coronal loops. In contrast to photospheric NLFFF codes, which calculate a magnetic field solution from the constraints of the transverse photospheric field, this new code uses coronal constraints instead, and this way provides important information on systematic errors of each magnetic field calculation method, as well as on the non-forcefreeness in the lower chromosphere. In this study we applied the COR-NLFFF code to active region NOAA 11158, during the time interval of 2011 Feb 12 to 17, which includes an X2.2 GOES-class flare plus 35 M and C-class flares. We calcuated the free magnetic energy with a 6-minute cadence over 5 days. We find good agreement between the two types of codes for the total nonpotential $E_N$ and potential energy $E_P$, but find up to a factor of 4 discrepancy in the free energy $E_{free}=E_N-E_P$, and up to a factor of 10 discrepancy in the decrease of the free energy $\\Delta E_{free}$ during flares. The coronal NLFFF code exhibits a larger time variability, and yields a decrease of free energy during the flare that is sufficient to satisfy the flare energy budget, while the photospheric NLFFF code shows much less time variability and an order of magnitude less free energy decrease during flares. The discrepancy may partly be due to the pre-processing of photospheric vector data, but more likely due to the non-forcefreeness in the lower chromosphere. We conclude that the coronal field cannot be correctly calculated based on photospheric data alone, but requires additional information on coronal loop geometries. ",
        "introduction": "The art of computing solar magnetic fields has evolved from routine potential field (PF) extrapolations to nonlinear force-free field (NLFFF) computations. The most important aspect concerning these two types of magnetic field models is the difference of magnetic energy content between them, the so-called {\\sl free energy}, which represents the maximum energy that can be dissipated during a magnetic instability that is driving solar flares, coronal mass ejections, or other eruptive events. Monitoring the magnetic evolution in active regions became an efficient diagnostic to measure the emergence of new magnetic flux on the solar surface, the magnetic storage, the injection of electric currents and helicity, the step-wise dissipation of nonpotential energy during flares, the energy conversion of magnetic energy during a magnetic reconnection process into particle acceleration, etc. The quantification of the free magnetic energy probably provides the strongest constraint to discriminate cause and consequence in the various energy conversion processes. It is therefore imperative to develop accurate and efficient tools to measure and monitor the evolution of magnetic energy in solar active regions. In this study we explore a novel method to calculate an approximate NLFFF solution using an automated loop tracing code, which aims to find force-free magnetic field solutions that are most consistent with the observed geometry of coronal loops, such as observed in high-resolution extreme-ultraviolet (EUV) images from the {\\sl Atmospheric Imaging Assembly (AIA)} (Lemen et al.~2012) onboard the {\\sl Solar Dynamics Observatory (SDO)} spacecraft (Pesnell et al.~2012).  NLFFF Modeling has been applied to active regions (Bobra et al.~2008; Su et al.~2009a; Inoue et al.~2011, 2013; Savcheva et al.~2012a), to coronal sigmoids (Savcheva et al.~2012b, 2012c; Inoue et al.~2012), to photospheric magnetic field changes (Liu et al.~2012), to magnetic flux emergence and energy build-up before flares (Su et al.~2009b; Li et al.~2007), to flares (Guo et al.~2008; Schrijver et al.~2008), to confined eruptions (Guo et al.~2010), to Coronal Mass Ejections (CMEs) (Su et al.~2011; Feng et al.~2013), as well as to MHD simulations (Savcheva et al.~2012a). Earlier studies compared theoretical models with observed images in a rather qualitative way, but recent studies go into quantitative comparisons of free energies and statistics of misalignment angles between modeled field lines and observed coronal loop geometries. Comparisons of up to 14 different NLFFF codes revealed substantial differences in the amount of calculated free magnetic energies, ranging from an (unphysical) minumun value $E_{N}/E_P=0.88$ below unity, to a maximum of $E_{N}/E_P=1.30$ for an X3.4-class flare (Schrijver et al.~2008). An evaluation of current NLFFF modeling was conducted in a study of the solar active region NOAA 10953 (2007 April 30) with a dozen of different NLFFF codes, which led to three major requirements for successful NLFFF modeling: (1) a sufficiently large field-of-view of the vector magnetic field data and the computation box, (2) accomodation of uncertainties in the boundary data, and (3) a realistic model of the transition from the non-forcefree photosphere to the force-free base of the corona (DeRosa et al.~2009). A pre-processing step to minimize the net force and torque in the photospheric boundary data was suggested (Wiegelmann et al.~2006, 2008; Wheatland and Regnier 2009), which improved the extrapolation above the chromosphere, but the field connectivity and free magnetic energy were not well recovered (Metcalf et al.~2008; Yamamoto et al.~2012; Jiang and Feng 2013). The latter issue is manifested in a major discrepancy between the computed magnetic field lines and stereoscopically observed geometries of coronal loops (Fig.~1), amounting to 3D misalignment angles of $\\approx 24^\\circ-44^\\circ$ (DeRosa et al.~2009). This mismatch angle can be reduced by forward-fitting of parameterized potential field models (Aschwanden and Sandman 2010; Sandman and Aschwanden 2011), by forward-fitting using a quasi-Grad-Rubin NLFFF method (Malanushenko et al.~2011, 2012), and by forward-fitting of analytical NLFFF approximations (Aschwanden et al.~2012, Aschwanden 2013a, 2013b, 2013c; Aschwanden and Malanushenko 2013). The latter method has been demonstrated to work equally well for both 3-dimensional (3D) loop coordinates, obtained with stereoscopy, and 2-dimensional (2D) loop projections (Aschwanden 2013c).  In this study we developed our forward-fitting code with analytical NLFFF approximations further by (1) implementing an automated loop tracing code for detection of coronal loops in multi-wavelength EUV images, which makes the manual or visual loop tracing unnecessary (Aschwanden 2010), (2) by optimization of the forward-fitting technique to 2D loop coordinates, which relinquishes 3D reconstruction with stereoscopy (Aschwanden 2013c), and (3) by improved detection and rejection of false loop structures, caused by CCD pixel bleeding, CCD saturation, and diffraction patterns from by the EUV telescope entrance filter. The content of this paper is a brief analytical description of the method (Section 2), a brief description of the numerical code (Section 3), and data analysis of NOAA active region 11158 observed with SDO/AIA and HMI during the 2011 February 12-17 flares (Section 4), discussion and comparison with previous studies on the same active region and X-class flare (Section 5), and conclusions (Section 6). ",
        "conclusions": "In this study we improved the accuracy and performance of the {\\sl Coronal Non-Linear Force-Free Field (COR-NLFFF)} forward-fitting code substantially by (1) implementing an automated loop tracing code for detection of coronal loops in multi-wavelength EUV images, which makes the manual or visual loop tracing unnecessary (Aschwanden 2010), and (2) by optimization of the forward-fitting technique to 2D loop coordinates, which relinquishes 3D reconstruction with stereoscopy (Aschwanden 2013c). In this study we applied the COR-NLFFF code to magnetic field modeling of AR 11158 during the time interval of 2011 Feb 12 to 17, which includes an X2.2 flare plus 35 M and C-class flares. We calcuated the free magnetic energy over the 5 days with a cadence of 6 minutes and compare the results with standard NLFFF calculations using the Wiegelmann code with a cadence of 12 minutes. We compare quantitatively the magnitude and evolution of the free energy during the GOES-detected flares, in particular the detected decreases of free energy before and after the flares with both types of codes. The standard NLFFF codes (Wiegelmann 2004; Jiang and Feng 2013) use extrapolation of photospheric vector magnetograph data (with or without pre-processing), while our forward-fitting COR-NLFFF code uses coronal loop constraints. We compare also calculations with the virial theorem. The results are the following:  \\begin{enumerate} \\item{The total nonpotential magnetic energy measured during the X2.2 flare on 2011 Febr 15, 02 UT, agrees well ($\\approx 15\\%$) between photospheric standard NLFFF codes and the coronal forward-fitting code, i.e., $E_{N} = (10.4\\pm1.4) \\times 10^{32}$ erg. However, the total energy determinied with the virial theorem is about a factor of two lower, i.e., $E_{N}=(4.97\\pm1.58) \\times 10^{32}$ erg, probably due to a different computation box.}  \\item{The potential energy measured during the X2.2 flare agrees also well ($\\approx 10\\%$) between the two types of codes, i.e., $E_P = (8.7\\pm0.8) \\times 10^{32}$ erg. The small difference, which is commensurable with the calibration of HMI and uncertainties among different NLFFF codes, corroborates the accuracy of the potential field modeling approach in terms of a a limited number of decomposed magnetic souces. The potential energy calculated with the virial theorem, $E_{P} = (2.43\\pm0.78) \\times 10^{32}$ erg, is a factor of 3 smaller, similar to the nonpotential energy, probably due to a different computation box.}  \\item{The free magnetic energy measured during the X2.2 flare, i.e., the difference between the nonpotential and potential energy, disagrees by a factor of about two between the two types of codes. For standard NLFFF codes we find a value of in the range of $E_{free}=(1.65-2.4) \\times 10^{32}$ erg, varying by about $15\\%$ with or without pre-processing, while we find a similar value of $E_{free}=1.0 \\times 10^{32}$ erg with the coronal forward-fitting COR-NLFFF code, so there is almost agreement between all codes in this case.}  \\item{The most critical quantity is the decrease of free magnetic energy before and after the X2.2 flare, which sets an upper limit on the energy budget of flares. We find a decrease in the free energy by $\\Delta E=E_{post}-E_{pre}=(0.37\\pm0.04) \\times 10^{32}$ erg with the standard photospheric NLFFF code (Sun et al.~2012a), which is about half of the decrease measured with the coronal forward-fitting COR-NLFFF code, i.e., $\\Delta E=E_{post}-E_{pre}=(0.62\\pm0.28) \\times 10^{32}$ erg, which is also consistent with the energy drop calculated with the virial theorem i.e., $\\Delta E=E_{post}-E_{pre}=(1.05\\pm1.04) \\times 10^{32}$ erg (Vemareddy et al.~2012), or $\\Delta E=E_{post}-E_{pre}=0.84 \\times 10^{32}$ erg (Tziotziou et al.~2013). The fact that photospheric NLFFF codes measure a higher amount of free energy than the coronal NLFFF codes, but a too low decrease during flares, could partly be due to a smoothing and time-averaging effect of the pre-preprocessing algorithm.}  \\item{The time evolution of the free energy, as calculated with a time resolution of 12 minutes over 5 days with both types, exhibits a systematic discrepancy in the amount of free energy obtained with the photospheric NLFFF versus the coronal NLFFF code. Since the free energy is dominated by the nonpotential transverse photospheric field, which corresponds to the azimuthal field component $B_{\\varphi}$ in helically twisted field lines, it is conceivable that the twist of field lines is not uniform along their length, but could be stronger twisted by about a factor of $\\sqrt{2}$ in the non-forcefree photosphere, and this way could explain why standard NLFFF codes produce a factor of $\\approx 2$ higher free energies than coronal-fitting NLFFF codes.}  \\item{Among the 36 GOES C,M, and X-class flares, we find a significant decrease in the free magnetic energy during the flares in 29 cases with the COR-NLFFF code. Similarly we find a significant decrease in many cases with the photospheric NLFFF code, however the energy drop is about a factor of $\\approx 10$ smaller for the photospheric NLFFF code than for the coronal (COR-NLFFF) code, and thus confirms statistically the same trend as we found for the X2.2 flare.}  \\item{The time evolution of the free energy in time steps of 12 minutes over 5 days exhibits about an order of magnitude stronger modulation depth for the coronal NLFFF code than for the photospheric NLFFF code, which is likely to be produced by the same effect that produces smaller energy decreases during the 36 flares, which could be partially be attributed to the smoothing effects of the pre-processing algorithm.}  \\end{enumerate}  This study represents the first comparative test case between photospheric and coronal NLFFF codes, which revealed discrepancies resulting from the inconsistency between the photospheric non-forcefree and the coronal force-free field. It appears that the so far existing preprocessing algorthm that supposedly optimizes the photospheric field into a force-free field, is not consistent with the coronal force-field extrapolated downward to a photospheric altitude level. Therefore we conclude that the coronal field cannot be retrieved from photospheric information alone, but requires explicit information from the magnetic field in the force-free corona. While the present COR-NLFFF code used here is based on an analytical approximation that is accurate two second order, more advanced NLFFF codes that include the geometry of coronal loops should be developed in future. One code in this direction, using a quasi-Grad-Rubin scheme is under current development (Anna Malanushenko, private communication), using visually traced loops as input. Of course, ultimate objectivity can only be achieved with automated loop recognition algorithms."
    },
    "1402/1402.5947_arXiv.txt": {
        "abstract": "The orbital dynamics of a test particle moving in the non-spherically symmetric field of a rotating oblate primary is impacted also by certain indirect, mixed effects arising from the interplay of the different Newtonian and post-Newtonian accelerations which induce known direct perturbations. We systematically calculate the indirect gravitoelectromagnetic shifts per orbit of the Keplerian orbital elements of the test particle arising from the crossing among the first even zonal harmonic $J_2$ of the central body and the post-Newtonian static and stationary components of its gravitational field. We also work out the Newtonian shifts per orbit of order $J_2^2$, and the direct post-Newtonian gravitoelectric  effects of order $J_2 c^{-2}$ arising from the equations of motion. In the case of both the indirect and direct gravitoelectric $J_2 c^{-2}$ shifts, our calculation holds for an arbitrary orientation of the symmetry axis of the central body. We yield numerical estimates of their relative magnitudes for systems ranging from Earth artificial satellites to stars orbiting supermassive black holes. \\textcolor{black}{As far as their measurability is concerned, highly elliptical orbital configurations are desirable.} ",
        "introduction": "Accurate tests of post-Newtonian gravity with either natural and artificial probes \\cite{Ginz59, 2010stsp.conf...25O, 2013CEJPh..11..531R, 2013NuPhS.243..172W} in a variety of astronomical and astrophysical scenarios as well as the long-term dynamics of hierarchical systems \\cite{2014PhRvD..89d4043W} require an ever increasing accuracy in modeling their orbital dynamics. In this respect, first-order perturbative calculations, yielding some of the most renown direct orbital consequences of the equations of motion \\cite{Sof89} like the Einstein perihelion precession \\cite{Ein15} and the Lense-Thirring effect \\cite{LT18},  should be complemented by second-order calculations. Indeed, they are able to capture also certain subtle consequences of the simultaneous presence of several terms in the equations of motion. They are indirect, mixed effects arising from the interplay among such terms which, in some cases, may have magnitudes comparable to some of the direct effects to the point that it has been recently stressed that they might be the subject of promising experimental investigations in a near future \\cite{2013CQGra..30s5011I}. While it is assumed that the orbital elements stay constant over one orbital revolution in calculating perturbatively the direct effects to the first order in some disturbing acceleration, accounting also for their instantaneous variations \\cite{2014PhRvD..89d4043W} yields either second-order and mixed effects if the acceleration is, actually, made of the sum of more than one term.  More specifically, let us consider a test particle moving in the non-spherically symmetric field of a rotating, oblate primary of mass $M$, mean equatorial radius $R$, quadrupole moment $J_2$ and angular momentum $\\bds S$: apart from the Newtonian monopole, the acceleration experienced by the particle (Section \\ref{pertu}) consists of the sum of a Newtonian term ${\\bds A}^{(J_2)}$ accounting for the primary's oblateness\\footnote{Here and throughout the paper, the other even zonal coefficients $J_{\\ell},\\ell=4,6,\\ldots$ of higher degree of the Newtonian multipolar expansion of the gravitational potential of the central body will be neglected.} $J_2$ and, to order $\\mathcal{O}\\ton{c^{-2}}$, of some static and stationary post-Newtonian terms ${\\bds A}^{(\\rm GE)},~{\\bds A}^{(\\rm GM)},~{\\bds A}^{(J_2~{\\rm GE})}$ yielding, to the first-order in each of them, direct effects like the gravitoelectric Schwarzschild-type  precession of the line of the apsides \\cite{Ein15}, the gravitomagnetic Lense-Thirring orbital precessions \\cite{LT18}, and some further orbital precessions proportional to $J_2 c^{-2}$ \\cite{1988CeMec..42...81S, 1991ercm.book.....B}. Actually, a perturbative calculation accounting also for the instantaneous shifts of the orbital elements during an orbital revolution is able to deliver, among other things, also mixed effects among such accelerations. It should be recalled that, in the perturbed 2-body Newtonian scenario, the 2nd order short term effects could be larger than the secular ones \\cite{2013osos.book.....X}. As a result, mixed orbital variations of order $\\mathcal{O}\\ton{J_2 c^{-2}}$, which are to be added to those directly arising from the equations of motion through ${\\bds A}^{(J_2~{\\rm GE})}$ \\cite{1988CeMec..42...81S, 1991ercm.book.....B}, occur. As recognized long ago, \\cite{Sof89, 1990CeMDA..47..205H}, they are of the same order of magnitude of the direct effects due to ${\\bds A}^{(J_2~{\\rm GE})}$. The gravitoelectric mixed effects were calculated in \\cite{2014PhRvD..89d4043W}, although they were not explicitly displayed. A calculation of them, based on Lie series and the Delaunay variables, can be found in \\cite{1990CeMDA..47..205H, arabo}. Moreover, there are also other mixed effects proportional to $J_2 S c^{-2}$ due to the interplay between the Newtonian quadrupolar field and the post-Newtonian gravitomagnetic one; to the best of our knowledge, they were never calculated in the literature. Direct orbital effects of order $\\mathcal{O}\\ton{J_2 S c^{-2}}$, calculated on the basis of a suitable multipolar expansion of the gravitomagnetic field of an axially symmetric source \\cite{1978PhRvD..18.1037T}, can be found in \\cite{1977PhRvD..16..946T}. For the direct post-Newtonian effects pertaining the precession of a gyroscope orbiting a rotating oblate body, see \\cite{1977PhRvD..16..946T, 1999GReGr..31.1837A, 2008arXiv0803.4133Z, 2014CQGra..31x5012P} As a by-product of such a calculation, also classical orbital shifts of order $\\mathcal{O}\\ton{J_2^2}$ can be obtained.  In this paper, we will analytically work out all the aforementioned effects (Section \\ref{misti}-Section \\ref{J2J2}) in a systematic and consistent way, outlined in Section \\ref{gene}, which, in principle, can be applied also to other disturbing accelerations, irrespectively of their physical origin. As far as both the mixed and the direct effects proportional to $J_2c^{-2}$ are concerned, we will calculate them for an arbitrary orientation of the spin axis of the central body. Our results, which are valid for generic orbital geometries of the test particle, represent the limits to which full two-body calculations must reduce in the point particle limit.  They can be used for sensitivity analyses involving different scenarios. Then, in Section \\ref{expe} we will numerically evaluate the relative strengths of such orbital shifts in various systems ranging from planetary artificial satellites \\cite{2007AcAau..61..932M, 2011AcAau..69..127P} to the stellar system orbiting the supermassive black hole (BH) located in Sgr A$^{\\ast}$ at the center of the Galaxy \\cite{2008ApJ...689.1044G}. Section \\ref{concludi} summarizes our findings. \\subsection*{Notations} Here,  basic notations and definitions used in the text are presented \\cite{1991ercm.book.....B, 2000Monte}. \\begin{description} \\item[] $G:$ Newtonian constant of gravitation \\item[] $c:$ speed of light in vacuum \\item[] $M:$ mass of the primary \\item[] $\\mu=GM:$ gravitational parameter of the primary \\item[] ${\\mathcal{R}}_s = 2\\mu c^{-2}: $ Schwarzschild radius of the primary \\item[] $R:$ mean equatorial radius of the primary \\item[] $J_2:$ dimensionless quadrupole mass moment of the primary \\item[] $S:$ angular momentum of the primary \\item[] ${\\kap}:$ unit vector of the spin axis of the primary \\item[] ${\\bds r}:$ radius vector of the test particle \\item[] $\\bds{\\hat{r}} = \\bds{r}/r:$ unit vector of the radius vector of the test particle \\item[] ${\\bds v}:$ velocity vector of the test particle \\item[] $\\bds k = \\bds r\\bds\\times \\bds v:$ orbital angular momentum per unit mass of the test particle \\item[] $\\bds{\\hat{k}}={\\bds k}/k:$ unit vector of the orbital angular momentum per unit mass of the test particle \\item[] $a:$  semimajor axis \\item[] $\\nk = \\sqrt{\\mu a^{-3}}:$   Keplerian mean motion \\item[] $P_{\\rm b} = 2\\pi \\nk^{-1}:$ Keplerian orbital period \\item[] $e:$  eccentricity \\item[] $p=a(1-e^2):$  semilatus rectum \\item[] $I:$  inclination of the orbital plane \\item[] $\\Om:$  longitude of the ascending node \\item[] $\\omega:$  argument of pericenter \\item[] $f:$  true anomaly \\item[] $u=\\omega + f:$  argument of latitude \\item[] $\\bds{\\hat{l}}:$ unit vector directed along the line of the nodes toward the ascending node \\item[] $\\bds{\\hat{m}}:$ unit vector directed transversely to the line of the nodes in the orbital plane \\item[] $\\bds{\\hat{P}}:$ unit vector directed along the line of the apsides toward the pericenter \\item[] $\\bds{\\hat{Q}}:$ unit vector directed transversely to the line of the apsides in the orbital plane \\item[] $\\bds A:$ disturbing acceleration \\item[] $A_R=\\bds A\\bds\\cdot{\\bds{\\hat{r}}}:$ radial component of $\\bds A$ \\item[] $A_T=\\bds A\\bds\\cdot\\ton{{\\bds{\\hat{k}}}\\bds\\times{\\bds{\\hat{r}}}}:$ transverse component of $\\bds A$ \\item[] $A_N=\\bds A\\bds\\cdot{\\bds{\\hat{k}}}:$ normal component of $\\bds A$ \\end{description} ",
        "conclusions": "\\lb{concludi} A first-order perturbative approach to particle dynamics in the post-Newtonian  field of a rotating oblate primary is not adequate to capture the full richness of the actual orbital motion due to the simultaneous contributions of several disturbing classical and relativistic accelerations ($J_2$, Schwarzschild, Lense-Thirring, etc.). Indeed, the very same fact that more than one  enter the equations of motion induces certain indirect, mixed orbital perturbations due to a mutual cross interaction in addition to second-order effects for each of them.  A consistent formalism able to reproduce such additional features of motion, which are not directly due to some new accelerations occurring in the equations of motion, is a second-order perturbative approach which we consistently outlined and applied to some known post-Keplerian accelerations of both Newtonian and post-Newtonian origin.  In particular, we considered the Newtonian acceleration induced by the oblateness $J_2$ of the central body, and the post-Newtonian gravitoelectromagnetic accelerations of order $\\mathcal{O}\\ton{c^{-2}}$ which, to first order, yield the well known Einstein and Lense-Thirring orbital precessions. We analytically calculated the indirect shifts per orbit of all the standard Keplerian orbital elements proportional to $J_2 c^{-2}$ and $J_2 S c^{-2}$. Our general approach is valid for an arbitrary orientation of the primary's spin axis $\\kap$. We also considered the Newtonian second-order effects in $J_2$. As far as the indirect gravitoelectric effects of order $\\mathcal{O}\\ton{J_2 c^{-2}}$ are concerned, they were added to the direct ones caused by the specific post-Newtonian acceleration proportional to $J_2 c^{-2}$ entering the equations of motion.  It turned out that the semilatus rectum $p$, the eccentricity $e$ and the inclination $I$ experience non-vanishing indirect shifts which are harmonic in the  argument of pericenter $\\omega$ entering their expressions as argument of trigonometric functions. The pericenter does not generally stay constant because of the direct perturbations of order $\\mathcal{O}\\ton{J_2}$ and $\\mathcal{O}\\ton{c^{-2}}$ that make it precess slowly. Instead, the node $\\Om$ and the pericenter itself undergo also indirect secular precessions because of some terms not containing explicitly $\\omega$. Our formulas, which are valid for a generic orbital geometry of the test particle, represent the limit to which full two-body formulas will have to reduce in the point particle limit.  Such indirect, mixed effects may play a role in realistic error budgets of accurate tests of post-Newtonian gravity and in the long-term evolutionary history of various astrophysical systems of interest. In principle, it is possible to design a dedicated satellite-based mission aimed to detect the effects of order $\\mathcal{O}\\ton{J_2c^{-2}}$ by looking at $p,e,I$. Indeed, in the Earth scenario, the dominant perigee precession is due to the Newtonian multipoles of the expansion of the terrestrial gravitational potential. Thus, a suitable orbital configuration, based on the concepts of critical inclination and frozen-perigee, may be adopted to suppress the largest part of the perigee precession as well as the long-term harmonic variations of the eccentricity and the inclination. In such a way, the shifts of order $\\mathcal{O}\\ton{J_2 c^{-2}}$ in $p,e,I$ would look like \\textcolor{black}{almost} secular trends over typical observational time spans some years long \\textcolor{black}{because of the very slow Schwarzschild-like gravitoelectric perigee advance}. As an example, a hypothetical terrestrial satellite orbiting at an altitude of $h=1~450$ km in an almost circular orbit inclined to the Earth's equator by an amount equal to the critical inclination able to suppress the Newtonian secular perigee precession due to $J_2$ as well as the long-term harmonic variations of the eccentricity and the inclination due to $J_3$, would experience an almost secular rate in $p$ as large as 51 cm yr$^{-1}$. Recent data analysis of the existing geodetic satellite LARES showed an accuracy in determining secular trends in the semimajor axis $a\\simeq p$ of the order of 14 cm yr$^{-1}$ over just one year. The eccentricity and the inclination would change at a rate of the order of $-11$ mas yr$^{-1}$ and $-3$ mas yr$^{-1}$, respectively. The nominal magnitude of the competing rates due to the atmospheric drag are much smaller."
    },
    "1402/1402.7167_arXiv.txt": {
        "abstract": "\\citet{wu09} suggested an empirical relation between the magnetic field strength along the line of sight (LOS) and the dispersion of Faraday rotation measure (RM) distribution in turbulent media with root-mean-square sonic Mach number $M_s \\simeq 1$. In this paper, we extend the work by incorporating the Mach number dependence. Media with $0.5 \\la M_s \\la 2$ are considered to cover the Mach number range of the warm ionized medium (WIM) of our Galaxy. Three-dimensional, magnetohydrodynamic isothermal turbulence simulations with solenoidal forcing are used. We suggest a new relation among the LOS magnetic field strength, the dispersion of RM distribution, and the Mach number, which approximately represents the relation for Alfv\\'enic perturbations. In addition, we suggest a relation between the Mach number and the dispersion of log-normal distribution of emission measure (EM), which is basically the relation for the Mach number and the density dispersion. The relations could be used for a quick and rough estimation of the LOS magnetic field strength in the turbulent WIM. ",
        "introduction": "The ionized gas in the Galaxy has been traditionally associated with bright regions surrounding hot stars, called Str$\\ddot{\\rm o}$mgren spheres or classical H II regions. Most H II regions, however, are found only at low Galactic latitudes with a scale height of $\\sim 40 - 70$ pc, which is much smaller than that of the diffuse ionized gas (DIG) or the warm ionized medium (WIM) \\citep[e.g.,][]{go01}. In fact, the classical H II regions contain only $\\sim 10 \\%$ of ionized hydrogen in the Galaxy, and the remaining $90 \\%$ resides in warm ($\\bar{T} \\sim 8000\\rm\\ K$) and diffuse ($\\bar{n}\\sim 0.03 \\ \\rm cm^{-3}$) regions. The WIM occupies approximately 20 - 30 \\% of the volume of $\\sim 2\\ \\rm kpc$-thick, plane-parallel layer of our Galaxy \\citep[e.g.,][]{reyn91,ha99,ga08,hi08}. The Wisconsin ${\\rm H}\\alpha$ Mapper (WHAM) northern sky survey has provided information on the distribution, kinematics, and other physical properties of the WIM \\citep{ha03}. The diffuse WIM is now recognized as one of the major components of the interstellar media (ISM) in both our Galaxy and external galaxies \\citep[e.g.,][]{ra90,reyn91,cr01,mv03} (see also \\citet{ha09} for a review).  The properties of the WIM have been revealed through observations of dispersion measure (DM) \\begin{equation} {\\rm DM} = \\int n_e ds \\end{equation} and emission measure (EM) \\begin{equation} {\\rm EM} = \\int n_e^2 ds, \\end{equation} where $n_e$ is the electron density and $s$ is the length along the line of sight (LOS). For instance, EM/DM and DM$^{2}$/EM give estimations of the electron density and the extent of the WIM, respectively \\citep[e.g.,][]{reyn91,be06,hi08}. In addition, the widths of probability density functions (PDFs) of DM and EM are known to be related to the root-mean-square (rms) sonic Mach number, $M_s$, of the WIM \\citep[e.g.,][]{bf08,hi08}.  The WIM is most likely to be in a state of turbulence, considering a very high Reynolds number inferred from observed velocities and expected viscosity \\citep{be99}. Evidences for turbulence include the log-normal distributions of 1) EMs from the H$\\alpha$ intensity in the Wisconsin H$\\alpha$ Mapper \\citep{ha99,ha03} survey \\citep{hi08}, 2) EMs from the H$\\alpha$ intensity in external galaxies \\citep[e.g.,][]{ta07,se09}, and 3) column densities of the DIG in our Galaxy \\citep{bf08}. The log-normality can be explained naturally by turbulent distribution of gas \\citep[e.g.,][]{vs94,es04,kri07,fe08,fe10}. The best evidence for turbulence comes from the power spectrum presented in \\citet{ar95}. It is a composite power spectrum of electron density collected from observations of velocity fluctuations of the interstellar gas, rotation measures (RMs), DMs, interstellar scintillations, and others. The spectrum covers a huge range of $\\sim 10^{10} - 10^{20}$ cm. The whole range of the spectrum is approximately fitted to the power spectrum of Kolmogorov turbulence with slope $-5/3$.  The WIM is permeated with magnetic fields. Measuring Faraday rotation of polarized radio emissions against background sources has been the most popular method for exploring magnetic fields there. The rotation of the plane of linearly polarized light is due to the birefringence of magneto-ionic medium. The angle of rotation is given by \\begin{equation} \\psi = {\\rm RM}\\ \\lambda^2, \\end{equation} where RM is the rotation measure defined by \\begin{equation} {\\rm RM} = 0.81 \\int n_e B_\\parallel ds. \\end{equation} Here, $\\lambda$ is the wavelength of polarized light and $B_{\\parallel}$ is the LOS magnetic field strength.  The units of RM, $n_e$, $B_\\parallel$ and $s$ are rad m$^{-2}$, cm$^{-3}$, $\\mu$G, and pc, respectively. RM/DM gives an estimation of the LOS magnetic field strength, weighted by the electron density, \\begin{equation} \\left< B_{\\parallel} \\right> = \\int n_{e} B_{\\parallel} ds\\ {\\Bigl /} \\int n_e ds. \\end{equation} The method, for instance, has been used to reconstruct the large-scale magnetic field in our Galaxy by many authors \\citep[e.g.,][]{han98,id99,fr01,han06,be07}.  \\citet{hav03,hav04} obtained the distributions of RMs along contiguous LOSs in the constellations of Auriga and Horologium in multi-frequency polarimetric observations with the Westerbork Synthesis Radio Telescope. While the peak in the frequency distribution of those RMs reflects the regular component of magnetic field, $B_0$, the spread should measure the turbulent component. This means that if a distribution of RM is observed, its spread provides another way to quantify the magnetic field in turbulent ionized media such as the WIM. Recently, using numerical simulations, \\citet{wu09} found that in turbulent media with $M_s \\simeq 1$, the width of the distribution of RM/$\\overline{\\rm RM}$ ($\\overline{\\rm RM}$ is the average value of RMs) is rather tightly related to the strength of the regular field along the LOS, $B_{0\\parallel}$. They suggested an empirical formula, which can be used to estimate $B_{0\\parallel}$ when the distribution of RM/$\\overline{\\rm RM}$ is available.  However, the Mach number is not necessarily unity in the WIM. The turbulent velocity dispersion has a range of values $v_{\\rm turb} \\sim 10 - 30$ km s$^{-1}$ \\citep{trh99}, and the temperature also has a range values $T \\sim 6000 - 12000$ K \\citep[e.g.,][]{re99}. So in the WIM, the turbulent flow motions, although not always having $M_s \\simeq 1$, are expected to be somewhere between mildly subsonic and mildly supersonic. As a matter of fact, several groups have suggested that the turbulent WIM has $M_s \\simeq 1 - 2.5$ \\citep[e.g.,][]{hi08,bu09,ga11,bu12}.  In this paper, we extend the work of \\citet{wu09} by incorporating the Mach number dependence. Specifically, we consider turbulent media with $M_s \\simeq 0.5$ and 2 along with $M_s \\simeq 1$. In Section 2, we outline our simulations. In Section 3, we present a new relation among $B_{0\\parallel}$, RM/$\\overline{\\rm RM}$, and $M_s$. We also suggest a relation between $M_s$ and the dispersion of log-normal distribution of EM, which can be used to estimate the Mach number. Summary follows in Section 4. ",
        "conclusions": "\\citet{wu09} found a relation between the magnetic field strength along the LOS, $B_{0\\parallel}$, and the FWHM of the frequency distribution of RM/$\\overline{\\rm RM}$, $W_{\\rm FWHM(RM/\\overline{RM})}$, for turbulent media of $M_s \\simeq 1$. But the Mach number in the WIM is not necessarily unity. In this paper, we incorporated the Mach number dependence in the relation. For it, we performed three-dimensional simulations of isothermal, compressible MHD turbulence for $M_s$ = 0.5, 1, 2 and $\\beta_0$ = 0.1, 1, 10. The parameters were chosen to cover the range of values expected in the WIM of our Galaxy.  - From the frequency distribution of RM/$\\overline{\\rm RM}$ which is well fitted with the Gaussian, we calculated $W_{\\rm FWHM(RM/\\overline{RM})}$. We suggest a relation among $B_{0\\parallel}$, $W_{\\rm FWHM(RM/\\overline{RM})}$, and $M_s$, which is shown in Equation (12).  - The frequency distribution of EM is well fitted with the log-normal distribution. We calculated the FWHM of the frequency distribution of $\\log_{10}\\rm (EM)$, $W_{\\rm FWHM[\\log_{10}({\\rm EM})]}$. We suggest a relation between $M_s$ and  $W_{\\rm FWHM[\\log_{10}({\\rm EM})]}$, which is shown in Equation (13).  The relation in Equation (12) would provide a handy way for a quick and rough estimation of $B_{0\\parallel}$ in the turbulent WIM regions where observations of RM distribution are available and $M_s$ is known. In the case that $M_s$ is unknown, the relation in Equation (13) could be used for an estimation of $M_s$, if observations of EM distribution are available."
    },
    "1402/1402.5448_arXiv.txt": {
        "abstract": "Electromagnetic metamaterials, made from arrangements of subwavelength sized structures, can be used to manipulate radiation. Designing metamaterials that have a positive refractive index along one axis and a negative refractive index along the orthogonal axis can result in birefringences, $\\Delta n>1$. The effect can be used to create wave plates with subwavelength thicknesses. Previous attempts at making wave plates in this way have resulted in very narrow usable bandwidths. In this paper, we use the Pancharatnam method to increase the usable bandwidth. A combination of Finite Element Method and Transmission Line models were used to optimise the final design. Experimental results are compared to the modelled data. ",
        "introduction": "Electromagnetic metamaterials, periodically arranged sub-wavelength structures, can be used to interact with radiation in ways generally unattainable with natural materials. In the form of metal mesh grids, they have been used since the 1960's by authors such as Ulrich~\\cite{ulrich_far-infrared_1967} to create polarisers and filters. The grids properties were based on the inductive and capacitive responses that radiation had to the grids geometries. In the past decade such grids have been rediscovered primarily for their use in producing negative refractive indices (NRI), and also for their ability to be designed with customised electromagnetic properties. This technology finds applications in the microwave and terahertz frequency ranges where the availability of suitable dielectric materials that can be used to manufacture optical components is very limited.  Most uses for NRI have focused on cases where isotropic behaviour is required e.g.\\ in the creation of ``perfect lenses'' \\cite{pendry_negative_2000}. However, polarisation dependent responses can be used to design birefringent metamaterials and hence wave plates. Wave plates (or retarders) are optical devices that alter the polarisation state of radiation transmitted through them. Half Wave Plates (HWPs) are a particular type of wave plate used to arbitrarily rotate the polarisation angle of linearly polarised radiation. They find application when linear polarisation modulation is required. In these cases, a HWP rotating at a frequency, $f$, is placed in front of a linearly polarised detector. In this setup, the intensity of linearly polarised radiation reaching the detector varies with a frequency of $4f$ whilst the intensity of unpolarised radiation is unaffected. This allows the detector to discriminate between radiation from polarised and unpolarised sources.  In this paper, NRI will be used in combination with positive refractive indices (PRI) to create a highly birefringent HWP. This idea has been applied in \\cite{weis_strongly_2009} where the results were very narrowband. In order in to increase the bandwidth in this work the Pancharatnam method \\cite{pancharatnam_achromatic_1955,pancharatnam_achromatic_1955-1} is employed for first time with NRI based wave plates. ",
        "conclusions": "In this paper, a highly birefringent metamaterial was created that was capable of producing a NRI and a PRI along orthogonal polarisation axes. This was used to create a HWP with a narrow usable fractional bandwidth defined as $\\Delta\\phi=\\SI{-180\\pm3}{\\degree}$ of \\SI{0.3}{\\percent}. Three of these HWPs were cascaded according to the Pancharatnam method to increase the usable bandwidth. TL based modelling of the Pancharatnam setup was able to increase this substantially, bringing the fractional bandwidth to \\SI{5.3}{\\percent}. Experimentally measured data taken with the manufactured Pancharatnam HWP showed a flattening of \\SI{3.1}{\\percent} in bandwidth between \\SIrange{88.9}{92.7}{\\GHz}. The reduced performance was due to the misalignment of the grids that made up the individual HWPs.  This work verifies the ability of the Pancharatnam method to create a broadening of the region where $\\Delta\\phi$ is flat even when the initial $\\Delta\\phi$ gradients of the constituent wave plates is very steep. Broader bandwidths would be achievable by cascading a greater number of HWPs, but this would come at the price of lower transmitted intensities."
    },
    "1402/1402.5162_arXiv.txt": {
        "abstract": "A sufficiently powerful astrophysical source with power law luminosity in time will give rise to a self-similar relativistic blast wave with a reverse shock traveling into the ejecta and a forward shock moving into the surrounding medium. Once energy injection ceases and the last energy is delivered to the shock front, the blast wave will transit into another self-similar stage depending only on the total amount of energy injected.  I describe the effect of limited duration energy injection into environments with density depending on radius as a power law, emphasizing optical / X-ray Gamma-ray Burst afterglows as applications. The blast wave during injection is treated analytically, the transition following last energy injection with one-dimensional simulations. Flux equations for synchrotron emission from the forward and reverse shock regions are provided. The reverse shock emission can easily dominate, especially with different magnetizations for both regions. Reverse shock emission is shown to support both the reported X-ray and optical correlations between afterglow plateau duration and end time flux, independently of the luminosity power law slope. The model is demonstrated by application to bursts 120521A and 090515, and can accommodate their steep post-plateau light curve slopes. ",
        "introduction": "\\label{introduction_section}  Some cataclysmic astrophysical events, such as the merger of neutron stars \\citep{Eichler1989, Paczynski1991} or the collapse of a very massive star \\citep{Woosley1993, Paczynski1998, MacFadyen1999}, can give rise to brief flashes of gamma rays (`gamma-ray bursts', or GRBs) that can be detected at cosmological distances. A common feature of GRB models is the launch of collimated relativistic ejecta that interact with the surrounding medium and produce an afterglow signal that can be detected from X-rays to radio as the blast wave decelerates.  The classical fireball model for the afterglow (e.g. \\citealt{Meszaros1993, Meszaros1997}) describes the relativistic hydrodynamical evolution after a mass and energy ($10^{48-51}$ erg, depending on the progenitor type) are injected effectively instantaneously into a small region. The fireball expands, accelerates and ultimately decelerates, by sweeping up external matter, with a self-similar relativistic fluid profile first described by \\cite{Blandford1976} (hereafter denoted BM76). This deceleration phase might be preceded by a brief phase where a reverse shock runs into the original ejecta (see e.g. \\citealt{Sari1995, Sari1997, Kobayashi1999}). The main observed emission component is synchrotron radiation from shock-accelerated electrons interacting with small-scale shock-generated magnetic fields, giving rise to a broadband signal that follows a power-law decay (or rise, at low observer frequencies) in time (e.g. \\citealt{Blandford1977}). Standard GRB afterglow theory has recently been reviewed by various authors, including \\cite{Piran2004, Meszaros2006, Granot2007, vanEerten2013proceedings}.  Since the launch of \\emph{Swift} \\citep{Gehrels2004}, early time plateau phases of shallow decay in afterglow X-ray light curves have been revealed to be far more common than originally expected. These plateaus are commonly attributed to prolonged injection of energy (see e.g. \\citealt{Nousek2006, ZhangBing2006}), which can take different forms, such as ejecta with a range of velocities catching up with the shock front (e.g. \\citealt{ReesMeszaros1998, Panaitescu1998, SariMeszaros2000}), long-term luminosity of the source (e.g. \\citealt{ZhangMeszaros2001}) or conversion of Poynting flux from the ejecta (e.g. \\citealt{Usov1992, Thompson1994, MeszarosRees1997Poynting, LyutikovBlandford2003}). When the blast wave is continously driven from the back, the impulsive energy injection scenario no longer correctly describes the fluid evolution and the blast wave will decelerate more slowly or not at all. Instead, for the duration of injection of energy, a more complex system of shocks will form similar to the brief reverse shock stage for massive ejecta, with a contact discontinuity separating the ejecta from the swept-up ambient medium and a reverse shock running into the ejecta, in addition to the forward shock running into the medium that is also present in the impulsive injection scenario.  The exact location and nature of emission during the plateau phase is still not fully resolved. The observed emission will be a mixture of forward and reverse shock contributions, and either can be dominant (see e.g. \\citealt{ZhangMeszaros2001, ZhangBing2006, Nousek2006, ButlerKocevski2007}). Treatments of outflows with a long-lived reverse shock have shown that it is possible to account for a significant part of the overall observed flux and light curve features by emission from the reverse shock region \\citep{Uhm2011, Uhm2012, Leventis2014}. Another strong indication that forward shock emission alone is insufficient to explain observations is the abrupt drop in luminosity that is sometimes seen at the end of the plateau phase (e.g. \\citealt{Troja2007, Rowlinson2013}).  In this study I consider the dynamics, evolution and emission from energy injection by a power law luminosity from the central source, which includes and generalizes the `thick-shell' case for massive ejecta. In the case of a sufficiently long-lived reverse shock that has become relativistic in the frame of the inflowing fluid, the full reverse shock / forward shock profile is self-similar and treated generally by BM76 (who omit only the reverse shock region density profile). Separate self-similar solutions exist in the literature (e.g. \\citealt{Nakamura2006, Nakayama2005}), but for our purpose, the radial flow with relativistic reverse shock described in BM76 is sufficient. Since the injection of energy still lasts only for a limited time, the transition from sustained to effectively impulsive injection (i.e. when the injection timescale becomes negligible again compared to the explosion duration) is discussed in detail, including a numerical hydrodynamics approach. Flux equations are derived for the emission that show the relative contributions from forward and reverse shock regions. These equations can be applied directly to observational X-ray and optical data. The reverse shock emission is found to often be important, as is demonstrated using `typical' long GRB afterglow parameters and short GRBs 090515 and 120521A.  In section \\ref{self-similar_section}, the self-similar fluid profile during energy injection from a power-law luminosity source is derived and placed in the context of the fireball model. In section \\ref{transition_section}, the transition after cessation of energy injection is discussed both analytically and numerically (using the \\textsc{ram} relativistic hydrodynamics code, \\citealt{ZhangWeiqun2006}). Flux equations and the role of the reverse shock region for the observed emission are treated in section \\ref{light_curve_section}, including an application to GRBs 090515 and 120521A. A general discussion follows in section \\ref{discussion_section}, a summary in section \\ref{summary_section}, and some technicalities are deferred to appendices.  In a separate study, \\cite{Leventis2014} also find an important role for the reverse shock emission in shaping the observed flux. They consider a homogeneous shell model for long-lived reverse shocks and assume a homogeneous circumburst medium, whereas this study consideres a generic power law medium, a full self-similar fluid profile and also examines the transition following energy injection using numerical simulations. Earlier work by \\cite{Uhm2012} considers the case where energy injection results from shells with a range of Lorentz factors catching up with the forward shock. There, the fluid profile is not self-similar but calculated in detail using the `mechanical model' approach to afterglow blast waves from \\cite{Beloborodov2006}. ",
        "conclusions": "\\label{discussion_section}  \\subsection{Charged particle acceleration and emission} \\label{cooling_subsection}  The detailed physics of charged particle acceleration and magnetic field generation in relativistic blast waves and turbulent flows are an incredibly complicated subject that is still poorly understood, and far beyond the scope of this paper. Instead, I have used the (commonly applied) simplifying parameters $\\epsilon_B$, $\\epsilon_e$ and $\\xi_N$, along with some implicit assumptions about their downstream evolution (see \\citealt{vanEerten2013proceedings} for details). Furthermore, a global approach was used to obtain a single cooling time for the entire plasma, rather than accounting for changing cooling times for fluid parcels as they advect away from acceleration sites. A more detailed treatment will shift the position of $\\nu_c$ within the RS and FS spectra (see also \\citealt{vanEerten2010offaxis}). Such a treatment, however, would need to address the question at which place(s) electrons are accelerated into a non-thermal distribution. Maybe the dominant process is Fermi shock acceleration at the FS and RS shock fronts. Local cooling could then, in principle, be calculated within the self-similar BM framework using the advection equation \\ref{soundwave_equation} from section \\ref{transition_section}. However, a (relativistic) FS-CD-RS system can easily be Rayleigh-Taylor unstable at the CD (see e.g. \\citealt{Duffell2013} for a recent demonstration). Particle acceleration can then also take place throughout the turbulent region, which will spread out through the RS and FS regions.  \\subsection{sideways spreading} \\label{sideways_spreading_subsection}  Once causal contact is established among all angles of the collimated blast wave, deviations from radial flow can be expected. This sideways spreading is partially responsible for what is observed as the `jet break' in the afterglow light curve (the other part being the edges of the outflow becoming visible). The onset of spreading can be calculated as the arrival time of a relativistic sound wave moving from edge to tip along a shock front decelerating according to $\\Gamma \\propto t^{-m/2}$. For nonzero $m$, this yields $\\theta_0 = 1 / [(3-k) \\Gamma]$ \\citep{vanEerten2013proceedings, MacFadyen2013}. For our typical ISM scenario with $q = 0$ and opening angle 0.1, this leads to $\\Gamma = 10$, which puts the onset of spreading well after $\\Gamma = 26.8$, when the last energy has been delivered to the shock front. In the typical wind scenario, $m = 0$, so we can not use $\\Gamma$ as a measure of the amount of time that has passed. In this case, $t = t_0 \\exp[4 \\theta_0 \\Gamma]$, for a sound wave departing from the edge at $t_0$. Since $t_0 \\sim R_0 / c$, causal contact along the shock front is therefore expected already early on during the injection phase.  Another way of assessing the importance of sideways spreading is by comparing injected energy and sideways energy loss, in terms of luminosities $L_{IN}$, $L_{OUT}$, respectively. Using the time derivative of the collimation-corrected amount of energy within the blast wave (approximated as a cylinder with homogeneous pressure set by the shock-jump conditions and width set by $\\chi = 1$, $\\chi_{RS}$) and sideways energy flux $F = \\sqrt{2} \\Gamma p c \\beta_\\perp'$ (where $\\beta_\\perp'$ the sideways velocity in the radially comoving frame, somewhere between sound speed $1 / \\sqrt{3}$ and light speed 1), we obtain \\begin{equation} \\frac{L_{IN}}{L_{OUT}} = \\frac{\\alpha (3 -k - m) \\sqrt{2}}{\\chi_{RS} \\beta_\\perp'} \\theta_0 \\Gamma, \\end{equation} where $\\alpha$ a term of order unity whose definition can be found in BM76. Since the other terms are also of order unity, the balance between energy flows is essentially set by $\\theta_0 \\Gamma$. The result is therefore similar to what is inferred from looking at sound waves along the shock front. For constant injection with $q = 0$, ISM blast waves remain radial until the jet break. Constant injection wind blast waves remain radial throughout the energy injection stage, except for a minor bow shock (if $\\Gamma > \\theta_0$), or start spreading immediately until $\\theta \\sim \\Gamma^{-1}$, which effectively puts a lower limit on how narrowly energy injection can be collimated when $m = 0$ (as is the case for a wind with $q=0$, $k =2$; note that the narrow jet treatment of GRB 090515 from the preceding section was done in the context of an ISM-type environment where $m=1$).  In a practical sense, this implies that the results obtained in this paper, both analytical and numerical, are approximately applicable even for blast waves that spread out at a later stage and no two-dimensional energy injection simulations are required. Simulation-based fit codes based on templates for impulsive energy injection blast waves \\citep{vanEerten2012boxfit, vanEerten2012scalings} can be extended to include an energy injection stage by attaching a series of one-dimensional simulation-based templates for various injection durations. As discussed previously in section \\ref{light_curve_section}, it is straightforward to extend scale invariance to include this stage.  What this does \\emph{not} mean, however, is that no jet break-like effect can be seen in the light curve during energy ejection. For nonzero $m$, there remains the effect of increasing relativistic beaming cone width even while the outflow remains radial. Given that afterglows are typically observed off-axis \\citep{Ryan2013, Ryan2014}, it is even likely that the onset of this type of jet break even occurs during the injection phase of the blast wave.  \\subsection{Energy injection by massive ejecta}  There exists an extensive literature on afterglow emission from massive ejecta (e.g. \\citealt{Sari1995, Kobayashi1999, Kobayashi2000, RamirezRuiz2002, Wu2003, Peng2005, Zou2005, Yi2013, Leventis2014}). Prolonged energy injection can be understood as a generalization of massive ejecta. In the latter case, the cold ejecta will stratify into a cold wind-type outflow (see also section \\ref{free_flowing_wind_subsection}) of width $\\Delta R \\propto R / \\Gamma^2$. As long as the RS has not fully crossed the ejecta, this situation is identical to a constant energy input (in the form of ejecta kinetic energy) with $q = 0$ (and in the `thick-shell' case, where the RS becomes relativistic in the frame of the ejecta before completing its crossing, the self-similar profiles described in this paper will also arise). Therefore, the flux equations derived in this paper will reduce to those presented e.g. by \\cite{Yi2013} if one takes $q = 0$. For ultrarelativistic ejecta, characteristic moments in the ejecta evolution like the completion of the reverse shock crossing and the deceleration radius are expected to occur in the observer frame on a timescale far smaller than the typically observed plateau end times of $10^{3-4}$ s. Less relativistic massive ejecta, such as the cocoon surrounding the collapsar jet, will have comparable timescales to the observed plateau durations. In that case the RS is likely to remain non-relativistic, so the full self-similar BM profile will not emerge. Nevertheless, equating the end-time of the plateau to the deceleration time and measuring the flux at this time will pick out the same point in the characteristic evolution across bursts, in the same way as measuring the flux at $T_{in}$ does.  \\subsection{The Lorentz factors of the outflow} \\label{Lorentz_factor_section}  Right now, the solutions presented in this work are based on a number of (connected) assumptions: the free-flowing relativistic wind is ultra-relativistic with Lorentz factor significantly exceeding that of the FS; The engine activity is long-lived, allowing for a relativistic RS to emerge; The Lorentz factor of the free-flowing wind remains fixed at $\\eta$, restricting the applicability of the solution to radii $> \\eta R_0$ and requiring an unchanging ratio between mass loss and luminosity of the engine. It should be noted that there are other ways of generalizing the standard fireball model of instantaneous massive ejecta, that do not necessarily lead to a relativistic RS. It is possible to relax the constraint that $\\eta$ remain fixed. If $\\eta$ changes over time, but the wind Lorentz factor at the RS remains larger than the FS Lorentz factor, the BM solution described in this paper still applies. The only difference then becomes that the time-dependence of $\\eta$ needs to be carried over to the density profile. If $\\eta \\propto t_{in}^s$, this can be implemented in the flux equations by replacing $\\eta \\to \\eta_{ref} (t_{obs} / t_{ref})^s$ (ignoring redshift effects; once again the factors $\\Gamma^2$ cancel in going from emission at the source to arrival at RS and from there to observer time). Relaxing the requirement that the engine remains active for a long time, quickly leads to a situation where the RS is not relativistic. At this point, the flux equations for the RS region (but not the FS region), quickly become qualitatively different. Under specific assumptions for the density profile of the ejecta and its acceleration behind the RS, self-similar solutions with Newtonian RS remain possible. An example can be found in the paper by \\cite{Nakamura2006} mentioned in the introduction. The flux equations for non-relativistic RS systems in general density profiles are provided by \\cite{Yi2013}."
    },
    "1402/1402.0437_arXiv.txt": {
        "abstract": "The AMS-02 detector is operating on the International Space Station since May 2011. More than 30 billion events have been collected by the instrument in the first two years of data taking. A precision measurement of the positron fraction and of the positron flux in primary cosmic rays up to 350 GeV, of the electron flux up to 500 GeV and of the combined electron plus positron flux up to 700 GeV are presented. The separate and combined electron and positron fluxes are preliminary and represent work in progress. Systematic uncertainties must still be investigated further. ",
        "introduction": "The Alpha Magnetic Spectrometer, AMS-02, is a general purpose high energy particle physics detector. It was installed on the International Space Station, ISS, on 19 May 2011 to conduct a unique long duration mission ($\\sim$20 years) of fundamental physics research in space.  \\begin{figure}[h!] \\centering \\includegraphics[width=0.4\\textwidth]{cospa13-1267-01} \\caption{\\emph{AMS-02 detector schematics in the event display of a 600 GeV electron.}} \\label{fig:detector} \\end{figure}  The layout of the AMS-02 detector~\\cite{bib:AMS02-PRL} is shown in Fig.~\\ref{fig:detector} presenting the event display of a 600 GeV electron recorded by AMS. It consists of nine planes of precision silicon tracker, a transition radiation detector (TRD), four planes of time of flight counters (TOF), a permanent magnet, an array of anticoincidence counters (ACC) surrounding the inner tracker, a ring imaging Cherenkov detector (RICH), and an electromagnetic calorimeter (ECAL). The TRD and the ECAL are separated by the Magnet and the Tracker. This ensures that secondary particles produced in the TRD and the upper TOF planes are swept away and do not enter into the ECAL.  The TRD is designed to use transition radiation to distinguish between electrons and protons, and {\\it dE/dx} to independently identify nuclei. It consists of 5248 proportional tubes of 6 mm diameter with a maximum length of 2 m arranged side by side in 16-tube modules. The 328 modules are mounted in 20 layers. Each layer is interleaved with a 20 mm thick fiber fleece radiator (LRP375) with a density of 0.06 $g/cm^3$. There are 12 layers of proportional tubes along the y axis located in the middle of the TRD and, along the x axis, four layers located on top and four on the bottom. The tubes are filled with a 90:10 Xe:CO$_{2}$ mixture. In order to differentiate between electrons and protons, signals from all the TRD layers are combined in a log-likelihood probability of the electron (TRD-LLe) or proton (TRD-LLp) hypothesis. The proton rejection power of the TRD estimator at 90\\,\\% e$^\\pm$ efficiency measured on orbit is $10^3$ to $10^4$, as shown in Fig.~\\ref{fig:rejection}a.  The ECAL consists of a multilayer sandwich of lead and scintillating fibers with an active area of 648$\\times$648 mm$^2$ and a thickness of 166.5 mm corresponding to 17 radiation lengths The calorimeter is composed of nine superlayers, each 18.5 mm thick. In each superlayer, the fibers run in one direction only. The 3D imaging capability of the detector is obtained by stacking alternate superlayers with fibers parallel to the x and y axes (five and four superlayers, respectively). The fibers are read out on one end by 1296 photosensors with a linearity of 1/10$^5$ per sensor. Signals from three super layers in y view (super layers 2,4,6) and in x view (super layers 1,3,5) are used in the trigger logic to select events with a shower in the calorimeter. The energy resolution of the ECAL is parametrized as a function of energy (in GeV) $\\sigma(E)/E=\\sqrt{(0.104)^2/E + (0.014)^2}$. In order to cleanly identify electrons and positrons, an ECAL estimator, based on a Boosted Decision Tree, BDT, algorithm~\\cite{bib06}, is constructed using the 3D shower shape in the ECAL. The proton rejection power of the ECAL estimator when combined with the energy-momentum matching requirement $E / p > 0.75$ reaches $\\sim$10,000 (see Fig.~\\ref{fig:rejection}b), as determined from the ISS data. The proton rejection power can be readily improved by tightening the selection criteria with reduced e$^\\pm$ efficiency.  \\begin{figure}[h!] \\centering \\includegraphics[width=0.9\\textwidth]{cospa13-1264-02} \\caption{\\emph{The measured proton rejection using the TRD} (a) \\emph{and the ECAL plus the Tracker} (b) \\emph{as a function of track momentum at 90\\,\\% selection efficiency for e$^\\pm$.}} \\label{fig:rejection} \\end{figure}  Two planes of TOF counters are located above and two planes below the magnet. Each plane contains eight or ten scintillating paddles. Each paddle is equipped with two or three photomultiplier tubes on each end for efficient detection of traversing particles.  The average time resolution of each counter has been measured to be 160 ps, and the overall velocity ($\\beta=v/c$) resolution of the system has been measured to be 4\\% for $\\beta\\simeq 1$ and Z=1 particles which also discriminates between downward and upward-going particles. The coincidence of signals from all four TOF planes provides the charged particle trigger.  The Tracker accurately determines the trajectory and absolute charge ($Z$) of cosmic rays by multiple measurements of the coordinates and energy loss. Coordinate resolution of each plane is measured to be better than 10$\\,\\mu$m in the bending direction and the charge resolution is $\\Delta Z \\simeq 0.06$ at $Z = 1$.  Together with the Magnet, the Tracker provides a Maximum Detectable Rigidity of 2 TV on average~\\cite{bib03}, over Tracker planes 1 to 9.  Specific  calibration procedures of all sub-detectors have been developed in order to guarantee the stability of the AMS-02 performances over time and no significant degradation of the apparatus has been observed during two years of operation in space~\\cite{bib:icrcTracker,bib:icrcTRD,bib:icrcTOF,bib:icrcECAL}. ",
        "conclusions": "The measurement of the electron spectrum with the AMS-02 detector has been performed at energies between 0.5 and 700 GeV and is reported in Fig.~\\ref{fig:measurement} (right top). The assessment of systematic uncertainties is currently being finalised. For this measurement, $\\sim$ 9 million electrons have been selected from more than 30 billion trigger collected in two years of operation in space. This represents $\\sim$ 10\\% of the expected AMS data sample.  The measured positron fraction is presented in Fig.~\\ref{fig:measurement} (left top) as a function of the reconstructed energy at the top of the AMS detector. The first 6.8 million primary positron and electron events collected with AMS on the ISS show: at energies $< 10\\,$GeV, a decrease in the positron fraction with increasing energy; a steady increase in the positron fraction from 10 to $\\sim$250$\\,$GeV; the slope of the positron fraction versus energy decreases by an order of magnitude from 20 to 250\\,GeV and no fine structure is observed.  Fig.~\\ref{fig:measurement} (left and right bottom) show the electron and positron fluxes. The electron flux measurement extends up to 500 GeV. Multiplied by $E^3$ it is rising up to 10 GeV and appears to be on a smooth, slowly falling curve above. The positron flux measurement extends up to 350~GeV. Multiplied by $E^3$  it is rising up to 10 GeV, from 10 to 30 GeV the spectrum is flat and above 30 GeV again rising as indicated by the black line in the figure. The spectral index and its dependence on energy is clearly different from the electron spectrum.  \\vspace*{0.3 cm} \\begin{figure}[h!] \\centering \\includegraphics[width=0.4\\textwidth]{cospa13-1264-05.pdf} \\quad \\includegraphics[width=0.49\\textwidth]{cospa13-1267-08} \\\\ \\includegraphics[width=0.45\\linewidth]{cospa13_1257_08top.pdf} \\quad \\includegraphics[width=0.45\\linewidth]{cospa13_1257_08bot.pdf} \\caption{(Left top) \\emph{The positron fraction compared with the most recent measurements from PAMELA~\\cite{bib15} and Fermi-LAT~\\cite{bib16}. The error bars for AMS are the quadratic sum of the statistical and systematic uncertainties and the horizontal positions are the centers of each bin.} (Right top) \\emph{AMS combined electrons+positron spectrum (red points) superimposed to recent measurements from references~\\cite{bib:ATIC,bib:PPB,bib:FERMI2010,bib:HESS2008,bib:HESS2009}.} (Left and right bottom) \\emph{AMS electron and positron spectra compared with the most recent measurements from PAMELA~\\cite{pamela} and HEAT~\\cite{heat}.}} \\label{fig:measurement} \\end{figure}"
    },
    "1402/1402.3328_arXiv.txt": {
        "abstract": "{High-mass stars are major players in the chemical and dynamical evolution of galaxies, and young massive clusters are the natural laboratories to study their evolution and their impact on star formation processes. Only in recent years have we become aware of the existence of numerous massive ($M_{{\\rm cl}}> 10^{4}\\:M_{\\odot}$) clusters in our Galaxy. Here I give a review, rather biased towards my own research interests, of the observational and theoretical efforts that have led to a description of their properties, and present an overview of the two (perhaps three) starburst regions known outside the Galactic Centre neighbourhood: the Scutum Complex, its putative counterpart on the far side of the Long Bar, and the starburst cluster Westerlund~1. } ",
        "introduction": " ",
        "conclusions": "In the past fifteen years, we have discovered almost fifteen open clusters more massive than $10^{4}\\:M_{\\odot}$ in the Milky Way. This has changed our perception of the Galactic cluster system. The concentration of these massive clusters towards the Sun suggests that many more are waiting to be discovered, in most cases hidden behind huge amounts of extinction.  Most of these clusters are relatively isolated, or surrounded by diffuse associations. Apart from the massive population near the Galactic Centre, the strongest concentration is the Scutum Complex, which extends over $\\sim400$~pc. At present, it is unclear if it represents a single entity, but the associations found surrounding Stephenson~2 and RSGC3 are likely the most massive ones known in the Milky Way, even if they are not connected. The most likely explanation for the existence of these large star-forming regions is triggering by the tip of the Long Bar. If so, a similar complex should exist on the other side of the Bar. Several massive clusters have been found in that direction. At least one of them, VdBH~222 is a cluster of red supergiants surrounded by a diffuse association, similarly to its near-side counterparts.  Even though Stephenson~2 is very massive, Westerlund~1 is likely to be the most massive cluster in the Milky Way. At an age between 4 and 6~Myr, its unique evolved population makes it the ideal laboratory to study high-mass star evolution. Wd~1 seems to have formed in isolation, and its radial velocity would place it at a distance difficult to reconcile with its observed population. These characteristics might be hinting at some violent dynamical process as the reason for its formation.    \\begin{acknowledgement} I would like to thank my collaborators in cluster research, Drs. Simon Clark, Carlos Gonz\\'alez-Fern\\'andez and Amparo Marco for many years of fruitful collaboration and comments on the manuscript. This research is partially supported by the Spanish Ministerio de Econom\\'{\\i}a y Competitividad under grants AYA2010-21697-C05-05 and AYA2012-39364-C02-02. \\end{acknowledgement}"
    },
    "1402/1402.3602_arXiv.txt": {
        "abstract": "This paper presents an absorption-line study of the multiphase circumgalactic medium (CGM) based on observations of \\lya, C\\,II, C\\,IV, Si\\,II, Si\\,III, and Si\\,IV absorption transitions in the vicinities of 195 galaxies at redshift $z<0.176$.  The galaxy sample is established based on a cross-comparison between public galaxy and QSO survey data and is characterized by a median redshift of $\\langle\\,z\\,\\rangle = 0.041 $, a median projected distance of $\\langle\\,d\\,\\rangle=362$ kpc to the sightline of the background QSO, and a median stellar mass of $\\log\\,(M_{\\rm star}/M_\\odot) = 9.7 \\pm 1.1$.  Comparing the absorber features identified in the QSO spectra with known galaxy properties has led to strong constraints for the CGM absorption properties at $z\\apl 0.176$.  First, abundant hydrogen gas is observed out to $d\\sim 500$ kpc, well beyond the dark matter halo radius $R_h$ of individual galaxies, with a mean covering fraction of $\\approx 60$\\%.  In contrast, no heavy elements are detected at $d\\apg 0.7\\,R_h$ from either low-mass dwarfs or high-mass galaxies.  The lack of detected heavy elements in low- and high-ionization states suggests that either there exists a chemical enrichment edge at $d\\approx 0.7\\,R_h$ or gaseous clumps giving rise to the observed absorption lines cannot survive at these large distances.  Considering all galaxies at $d>R_h$ leads to a strict {\\it upper limit} for the covering fraction of heavy elements of $\\approx 3$\\% (at a 95\\% confidence level) over $d=(1-9)\\,R_h$. At $d<R_h$, differential covering fraction between low- and high-ionization gas is observed, suggesting that the CGM becomes progressively more ionized from $d<0.3\\,R_h$ to larger distances. Comparing CGM absorption observations at low and high redshifts shows that at a fixed-fraction of $R_h$ the CGM exhibits stronger mean absorption at $z=2.2$ than at $z\\sim 0$ and that the distinction is most pronounced in low-ionization species traced by C\\,II and Si\\,II absorption lines.  We discuss possible pseudo-evolution of the CGM as a result of misrepresentation of halo radius, and present a brief discussion on the implications of these findings. ",
        "introduction": "Circumgalactic space is thought to contain a vast amount of baryons (e.g.\\ Spitzer 1956; Fukugita 2004) that are both multiphase and dynamic (e.g.\\ Mo \\& Miralda-Escud\\'e 1996, Maller \\& Bullock 2004).  In principle, this circumgalactic gas provides the fuel necessary for sustaining star formation in galaxies, while at the same time it is also being replenished with both newly accreted (and presumably chemically pristine) intergalactic gas and chemically enriched materials either through gas stripping due to satellite interaction or through starburst driven winds.  Therefore, circumgalactic space provides a critical laboratory for studying the baryon cycle that regulates star formation and galaxy growth.    Over the past two decades, absorption spectroscopy of background QSOs have provided an effective tool for probing the otherwise unseen, diffuse gas in halos around distant galaxies.  While each QSO provides a one-dimensional mapping through each single halo, observing an ensemble of close galaxy and QSO pairs leads to statistical measurements of mean halo profiles (i.e.\\ absorption strength as a function of projected distance) averaged over the entire galaxy sample.  Commonly seen transitions in circumgalactic space include hydrogen \\lya\\,$\\lambda\\,1215$ (e.g.\\ Lanzetta \\etal\\ 1995; Chen \\etal\\ 1998, 2001a; Tripp \\etal\\ 1998; Wakker \\& Savage 2009; Steidel \\etal\\ 2010; Prochaska \\etal\\ 2011; Thom \\etal\\ 2012; Stocke \\etal\\ 2013; Rudie \\etal\\ 2013), the C\\,IV\\,$\\lambda\\lambda\\,1548, 1550$ doublet (e.g.\\ Chen \\etal\\ 2001b; Adelberger \\etal\\ 2005; Steidel \\etal\\ 2010), the Mg\\,II\\,$\\lambda$$\\lambda\\,2796,2803$ doublet (e.g.\\ Bowen \\etal\\ 1995; Chen \\& Tinker 2008; Kacprzak \\etal\\ 2008; Barton \\& Cooke 2009; Chen \\etal\\ 2010a,b; Gauthier \\etal\\ 2010; Bordoloi \\etal\\ 2011; Werk \\etal\\ 2013), and the O\\,VI\\,$\\lambda\\lambda\\,1031,1036$ doublet (Chen \\& Mulchaey 2009; Wakker \\& Savage 2009; Prochaska \\etal\\ 2011; Tumlinson \\etal\\ 2011; Johnson \\etal\\ 2013).  A general finding from different studies is that most galaxies are surrounded by chemically enriched gas out to $\\approx 100-300$ kpc in projected distance.  Competing scenarios for explaining the presence of heavy elements at $> 50$ kpc from star-forming regions include super galactic winds (e.g.\\ Murray \\etal\\ 2011; Booth \\etal\\ 2013) and tidally disrupted satellites (e.g.\\ Wang 1993; Gauthier 2013).    A particularly interesting absorption feature to adopt for probing the baryon content of galactic halos is the \\civ\\ doublet.  These doublet transitions are strong and their rest-frame wavelengths, 1548.20 and 1550.77 \\AA, enable uniform surveys of chemically enriched gas in the optical window over a broad redshift range from redshift $z\\approx 1$ to $z\\apg 5$.  \\civ\\ absorbers are found to originate primarily in photoionized gas of temperature $T\\approx 4 \\times 10^4$ K with some contribution from shock heated gas in galaxies and galactic halos (e.g.\\ Rauch \\etal\\ 1997; Boksenberg \\etal\\ 2003; Simcoe \\etal\\ 2004). Therefore, \\civ\\ absorption transitions provide an effective tracer of chemically enriched warm gas in and around star-forming regions.  A number of surveys have been carried out to characterize the statistical properties of \\civ\\ absorbers found along random QSO sightlines (e.g.\\ Songaila 2001; Boksenberg \\etal\\ 2003; Simcoe 2011; Cooksey \\etal\\ 2013).  While the shape of the \\civ\\ equivalent width frequency distribution function appears to remain the same, the total cosmic mass density in C$^{3+}$ ions is found to increase from $z=4$ to $z\\approx 1.5$ by a factor of 2 (e.g.\\ Cooksey \\etal\\ 2013).  The increasing \\civ\\ mass density together with an increasing background radiation intensity with decreasing redshift (e.g. Haardt \\& Madau 2012) indicate an increasing chemical enrichment level with time in the gas traced by the \\civ\\ absorption transitions (e.g.\\ Oppenheimer \\& Dav\\'e 2008).   At the same time, only three studies have been carried out to characterize extended \\civ\\ gaseous halos around galaxies, two at $z\\apl 0.5$\\footnote{Though we have recently learned that Bordoloi \\etal\\ (2014) has also been conducting a survey of \\civ\\ absorbers in halos around low-redshift galaxies.} and one at $z\\approx 2.2$. A clear understanding has yet to be established.  At $z\\approx 0.4$, Chen \\etal\\ (2001b) studied the incidence of \\civ\\ absorbing gas at projected distances $d\\apl 300$ kpc from a sample of 50 galaxies. These authors reported the presence of a distinct boundary at a $B$-band luminosity-normalized projected distance of $\\hat{d}\\equiv d\\times(L_B/L_B^*)^{-0.5}\\approx 160$ kpc, beyond which no \\civ\\ absorbers are detected.  The lack of \\civ\\ absorption is based on a subsample of 35 galaxies that occur at $\\hat{d}> 160$ kpc from their background QSO sightlines.  In contrast, 14 of the 15 galaxies at $\\hat{d}<160$ kpc have an associated \\civ\\ absorber, although the rest-frame absorption equivalent width, \\ewr, exhibits a large scatter between individual galaxies.  At $z\\approx 2.2$, Steidel \\etal\\ (2010) examined the mean spatial absorption profiles of 512 galaxies based on stacked spectra of background galaxies that occur at $d\\apl 125$ kpc. These authors found that the mean \\civ\\ absorption strength declines rapidly at projected distances of $\\approx 50-100$ kpc.  Comparing the observations of Chen \\etal\\ (2001b) and Steidel \\etal\\ (2010) yielded little distinction in the \\civ-traced circumgalactic medium (CGM) at low and high redshifts (Chen 2012).  Both the luminosity-normalized spatial extent and mean absorption equivalent width of the CGM around galaxies of comparable mass (but with very different on-going star formation rate) have changed little over the redshift interval $z=0.4-2.2$, although there exists a large scatter in \\ewr\\ in the low-redshift study.  Nevertheless, a lack of variation in the spatial profile of the chemically enriched CGM between two distinct epochs poses a serious challenge to the theoretical models of gas flows around galaxies (e.g.\\ Hummels \\etal\\ 2013; Ford \\etal\\ 2013a,b).  On the other hand, Borthakur \\etal\\ (2013) targeted a sample of 20 galaxies at $z<0.2$ and were able to constrain the incidence of \\civ\\ absorbing gas for 17 of these galaxies.  Of the 17 galaxies studied, eight occur at $\\hat{d}<160$ kpc and nine at larger distances.  These authors detected \\civ\\ absorbers for three of the eight galaxies at $\\hat{d}<160$ kpc and three of the nine galaxies at $\\hat{d}>160$ kpc in their sample.  The finding of a flat rate of \\civ\\ incidence with increasing projected distance is intriguing and appears to be in stark contrast to the earlier finding of Chen \\etal\\ (2001b).  While Borthakur \\etal\\ attributed the detected \\civ\\ absorbers to starburst activity in the associated galaxies, it is not straightforward to reconcile the discrepant trend found for extended \\civ\\ halos between the Chen \\etal\\ and Borthakur \\etal\\ samples.%    To address the discrepant trend found for extended \\civ\\ halos around galaxies at low redshifts and to improve the empirical understanding of how the chemically enriched CGM evolves with time, we have searched public archives to find spectroscopically identified galaxies at small projected distance to a UV bright, background QSO for which high-quality echelle/echellette spectra are already available in the {\\it Hubble Space Telescope} (HST) data archive.  Our search has yielded the first large sample of $\\sim 300$ close QSO and galaxy pairs that enables a systematic study of extended \\civ\\ halos over the projected distance interval of $d = 0-500$ kpc.  As described below, this unique sample allows us to place unprecedented limits on the extent of chemical enrichment around low-redshift galaxies based on observations of not only the \\civ\\ absorption doublets but also a whole host of ionic transitions from low- and intermediate-ionization states.  Recently, Tumlinson \\etal\\ (2013) carried out a large program to use the Cosmic Origins Spectrograph (COS; Green et al.\\ 2012) on board HST to study the CGM at $d<150$ kpc of 44 $L_*$ galaxies at $z=0.15-0.35$. This program, known as the ``COS-Halos'' survey, was designed to map the multi-phase CGM using O\\,VI and other metal-line diagnostics (Tumlinson \\etal\\ 2011; Werk \\etal\\ 2013).  The low-redshift cut at $z=0.15$ was dictated by the requirement of detecting the O\\,VI doublet using COS and the FUV channel.  However, the spectral coverage of the COS FUV gratings misses the C\\,IV doublet at $z\\apg 0.15$, and therefore \\civ\\ transitions are not included in the COS-Halos analysis (e.g.\\ Werk \\etal\\ 2013).  Our study differs from the COS-Halos program in two fundamental aspects.  First, our search is designed to probe the \\civ\\ absorption in galactic halos for comparisons with high-redshift studies. Second, because fainter galaxies are also more numerous, it is more likely to find a faint foreground galaxy near a QSO sightline than a luminous galaxy (see the discussion is \\S\\ 2).  As a result, low-luminosity ($<0.1\\,L_*$) and low-mass ($<0.1\\,M_{\\rm star}^*$) dwarf galaxies contribute to a large fraction ($\\approx 50$\\%) of our random QSO and galaxy pair sample.  Therefore, {\\it our study complements the COS-Halos effort both in terms of the mass regime of the galaxies and in terms of the ionization state of the CGM}.   Here we report initial findings for the low-redshift CGM based on a sample of 195 galaxies for which constraints on the stellar mass are available.  The available stellar mass estimate of each galaxy allows us to infer the total mass of the dark matter in which the galaxy resides (e.g.\\ Behroozi et al.\\ 2013; Kravtsov et al.\\ 2014). This in turn allows us to assess the intrinsic size differences between different galaxies based on the estimated dark matter halo radius from a standard halo model.  While the CGM is expected to be regulated by various complicated physical processes, including gravitational motion and star formation feedback, addressing the effects of all these processes at once is challenging.  As a first step toward establishing a clear understanding of the processes that drive the observed CGM properties, we have designed our study here to account for the intrinsic size differences (driven by the gravitational potential) of individual galactic halos by normalizing the observed projected distances with the underlying halo radius.  We expect that the results will facilitate future studies that consider the effect of additional processes such as star formation feedback.  This paper is organized as follows.  In Section 2, we describe the selection criteria for establishing the close galaxy--QSO pair sample and summarize the properties of the galaxy members in the pair sample. In Section 3, we describe the data reduction and analysis procedures of the absorption spectra of the QSOs.  In Section 4, we examine the correlation between the strengths of different ionic transitions and galaxy properties, and compare the results of our study with previous findings.  Finally, we discuss the implication of our results in Section 5.  Throughout the paper, we adopt the standard $\\Lambda$ cosmology, $\\Omega_\\Lambda = 0.7$, $\\Omega_M = 0.3$ with a Hubble constant $H_0 = 70\\,\\kms\\,{\\rm Mpc}^{-1}$. ",
        "conclusions": "Based on absorption spectroscopy carried out in the vicinities of 195 galaxies at $z\\apl 0.17$, we have obtained strong constraints for the absorption properties of the CGM at low redshift.  We observe distinct spatial absorption profiles between H\\,I and heavy elements, with extended \\lya\\ absorbing gas found at distances far beyond the halo radius $R_h$ but no heavy ions detected at $d\\apg 0.7\\,R_h$.  Here we discuss the implications of our study.  \\subsection{Observed Absence of Heavy Elements at $d\\apg 0.7\\,R_h$ from Star-forming Regions}  A particularly interesting finding from our absorption-line search is that while between 60\\% and 90\\% of the galaxies at $d\\apg 0.7\\,R_h$ show associated \\lya\\ absorbers in the QSO spectra, all but two of these \\lya\\ absorbers exhibit associated heavy ions to sensitive upper limits.  The lack of detected heavy ions applies to both low-ionization species probed by the C\\,II and Si\\,II absorption transitions and high-ionization species probed by Si\\,III, Si\\,IV, and C\\,IV.  It implies that either the chemical enrichment in galactic halos has a finite edge at $d\\approx 0.7\\,R_h$ or gaseous clumps giving rise to the observed absorption lines cannot survive at these large distances.  In addition, the observed absence of heavy ions at $d\\apg 0.7\\,R_h$ is found around galaxies that span over four orders of magnitudes in $M_{\\rm star}$ (Figure 3).  While the $R_h$-normalized radial profile in principle accounts for the size-scaling between galaxies of different mass, it does not include possible variations of halo gas properties as a result of the star formation history of the galaxies. Specifically, star formation rate is known to correlate strongly with galaxy mass (e.g.\\ Figure 4) and star formation feedback is expected to substantially alter the physical properties of halo gas (e.g.\\ Agertz \\& Kravtsov 2014).  At the same time, high-mass galaxies of $M_{\\rm star}\\sim 10^{11}\\,M_\\odot$ are often found in quiescent mode with little on-going star formation (e.g.\\ Muzzin \\etal\\ 2013), which would naturally result in a reduced absorption strength in halos around these massive galaxies (e.g.\\ Gauthier \\etal\\ 2010).  To investigate possible variation in the $R_h$-normalized radial absorption profile of the CGM, we divide the galaxy sample into low- and high-mass subsamples.  Figure 12 shows the \\civ\\ absorption properties for low-mass galaxies of $M_{\\rm star}\\apl 0.1\\,M_{\\rm star}^*$ (filled and open squares) and high-mass galaxies of $M_{\\rm star}\\apg 0.1\\,M_{\\rm star}^*$ (filled and open circles), where $M_{\\rm star}^*=5\\times 10^{10}\\,{\\rm M}_\\odot$ (Baldry \\etal\\ 2012; Muzzin \\etal\\ 2013).  Both low- and hig-mass galaxies exhibit a modest mean covering fraction of strong ($W_r(1548)>0.1$ \\AA) \\civ\\ absorbers $\\mkap\\approx 30-50$\\% at $d< 0.7\\,R_h$, and no associated \\civ\\ absorbers at $d>R_h$.  The sample is too small to make a strong distinction, beyond $R_h$-normalization, in the extent of chemical enrichment between low- and high-mass halos in our sample.  \\begin{figure} \\centering \\includegraphics[scale =0.425]{mass_civ_rh_sb} \\caption{$R_h$-normalized radial profiles of \\civ\\ absorbing gas for low-mass galaxies of $M_{\\rm star}<0.1\\,M_{\\rm star}^*$ (filled and open squares) and high-mass galaxies of $M_{\\rm star}>0.1\\,M_{\\rm star}^*$ (filled and open circles), where $M_{\\rm star}^*=5\\times 10^{10}\\,{\\rm M}_\\odot$ (Baldry \\etal\\ 2012; Muzzin \\etal\\ 2013). Constraints from stacked spectra of the full sample are included for reference (pentagon points).  Both low- and hig-mass galaxies exhibit a modest mean covering fraction of strong ($W_r(1548)>0.1$ \\AA) \\civ\\ absorbing gas $\\mkap\\approx 30-50$\\% at $d< 0.7\\,R_h$, and no associated \\civ\\ absorbers at $d>R_h$.  No significant distinction, beyond $R_h$-normalization, can be made in the extent of chemical enrichment around low- and high-mass galaxies.  In contrast, six of the 17 galaxies in the Borthakur \\etal\\ (2013) sample exhibit strong \\civ\\ absorption out to $\\sim R_h$ (star symbols).  While the constraints for the 11 non-detections are weak, the detected \\civ\\ absorbers around six galaxies (four starburst and two passive galaxies) are all significantly stronger than the mean values observed in our sample. } \\end{figure}  A lack of heavy ions outside of the halos where dwarf galaxies reside provides important constraints for the amount of heavy elements that can escape these low-mass halos.  Starburst driven outflows are thought to be particularly effective in removing baryons from low-mass halos because of a shallow potential well (e.g.\\ Larson 1974; Martin 1999).  Energetic outflows are also expected to suppress continuing growth of low-mass galaxies due to a reduced accretion rate (e.g.\\ Scannapieco \\etal\\ 2001).  These effects together can explain known scaling relations such as the mass-metallicity relation found in dwarf galaxies (e.g.\\ Dekel \\& Woo 2003).  However, model that predict enrichment levels to $>0.1$ solar metallicity are inconsistent with the observed limits we obtain at $d>0.7\\,R_h$ from low-mass dwarf galaxies at $z\\sim 0$.  This places an important constraint on feedback models that would enrich the CGM to much higher metallicity at these distances.  Our finding of a lack of detected heavy elements beyond $d\\approx 0.7\\,R_h$ appears to be discrepant from the observations of Borthakur \\etal\\ (2013), who reported detections of strong \\civ\\ out to $d\\sim R_h$ from massive galaxies of $\\log\\,(M_{\\rm star}/M_\\odot)\\approx 10.1$ (star symbols in Figure 12).  A possible explanation is to attribute the observed \\civ\\ absorbers to outflows driven by on-going star-formation activity, which indeed was part of the selection criterion of the Borthakur \\etal\\ sample.  However, two of the galaxies with associated strong \\civ\\ absorbers are classified as quiescent galaxies with little on-going star formation since recent past.  An alternative scenario different from starburst driven outflows in explaining the presence of heavy elements at large distances is through gas stripping in an overdense group environment (e.g.\\ Kawata \\& Mulchaey 2008; McCarthy \\etal\\ 2008).  Such scenario can be tested with deep galaxy survey data of fields around these strong metal-line absorbers (e.g.\\ Gauthier 2013).  To improve upon the detection limits of heavy ions in \\lya\\ absorbing gas, we repeat the stacking exercise discussed in \\S\\ 4.4 but consider only galaxies at $d\\apg R_h$ which exhibit associated \\lya\\ absorbers. The goal of this exercise is to investigate whether some constraints can be obtained on the chemical enrichment level in known gaseous clouds revealed by \\lya\\ absorption.  The absorption-line measurements and constraints in stacked spectra of \\lya\\ absorbers are summarized in Table 8.  We divide the galaxies at $d\\apg R_h$ into three subsamples, each covering different $R_h$-normalized projected distance intervals.  The goal is to maximize the $S/N$ of the stacked spectra, while searching for possible trend of chemical enrichment with distance.  In this particular exercise, we also search for the presence of the N\\,V\\,$\\lambda\\lambda\\,1238, 1242$ doublet features.  Similar to the O\\,VI absorption doublet, the N\\,V doublet is seen in highly ionized gas and is covered in the STIS and COS FUV spectral range. Observations of the N\\,V absorption doublet together with C\\,IV and other low-ionization transitions, therefore, provides a measure of the ionization state of the gas (see for example Figure 32 of Thom \\& Chen 2008).  The results of our search remains the same; namely no ionic transitions are detected in \\lya\\ absorbers at $d\\apg R_h$ from galaxies in our sample.  Adopting the absorption constraints for \\lya\\ and C\\,IV from Table 8 and assuming typical Doppler parameters, $b({\\rm H\\,I})=30$ \\kms\\ (e.g.\\ Tilton \\etal\\ 2012) and $b({\\rm C\\,IV})=7$ \\kms\\ (e.g.\\ Boksenberg \\etal\\ 2003), we infer a column density ratio of $\\log\\,N({\\rm C\\,IV})/N({\\rm H\\,I})<-1.3$.  In contrast, cosmological simulations that incorporate strong outflows to reproduce statistical properties of the general galaxy population in the low-redshift universe predict a column density ratio of $\\log\\,N({\\rm C\\,IV})/N({\\rm H\\,I})\\sim -0.6$ at the virial radius of galactic halos of $10^{11-12}\\,{\\rm M}_\\odot$ (Ford \\etal\\ 2013b). For photo-ionized gas, this column density ratio implies that the chemical enrichment level of \\lya\\ absorbing gas at $d\\apg R_h$ cannot exceed 0.1 solar metallicity.  \\subsection{The Origin of a Chemical Enrichment Edge at $\\sim 0.7\\,R_h$}  The distinct boundary between the presence and absence of metal-line absorbers at $d\\approx 0.7\\,R_h$ is seen in all ionic transitions that have been studied, including the observations of Si\\,II, Si\\,III, Si\\,IV, C\\,II, and C\\,IV reported here and those of Mg\\,II absorbers around $\\langle\\,z\\,\\rangle=0.24$ galaxies (Chen \\etal\\ 2010a).  It is, however, not clear whether such absorption boundary exists for highly ionized gas probed by the O\\,VI absorption doublet.  Existing observations are primarily limited to inner galactic halos at $d\\apl 0.5\\,R_h$ (e.g.\\ Tumlinson \\etal\\ 2011, 2013).  Other studies that extend to larger radii have only small samples of ``isolated'' galaxies (e.g.\\ Chen \\& Mulchaey 2009; Prochaska \\etal\\ 2011) and therefore available constraints are uncertain (see Johnson \\etal\\ 2014, in preparation).  A natural explanation for the chemical enrichment edge at $d\\approx 0.7\\,R_h$ is the turn-around radius $r_{\\rm turn}$ of starburst driven outflows, similar to a halo fountain phenomenon (cf.\\ Bregman 1980). We note that as chemically-enriched outflowing material moves into low-density regions and mixes in with metal-poor gas, it is expected that the outflowing material would be diluted and the observed absorption strength would gradually decline.  However, to produce the sharp downturn at $d\\approx 0.7\\,R_h$ would require a sharp decline in metallicity, which essentally indicates a finite edge of chemical enrichment and motivates the consideration of a turn-around radius of starburst driven outflows.  In this scenario, we can infer the launch speed as a function of launch radius of the outflows in individual halos.  Adopting a Navarro-Frenk-White (NFW; Navarro \\etal\\ 1997) density profile for the underlying dark matter halos, we estimate the launch velocity $v_{\\rm out}$ at a launch radius $r_i$ following $v_{\\rm out}(r_i)=\\sqrt{2\\,[\\Phi(r_{\\rm turn})-\\Phi(r_i)]}$, where $\\Phi$ is the gravitational potential and is related to radius $r$ according to $\\Phi(r)=-(GM_s/r)\\times\\ln\\,(1+r/r_s)$.  For an NFW profile, $r_s$ is the scale radius and $M_s$ is the total mass enclosed within $r_s$ and is related to the halo mass according to $M_h=M_s\\,[\\ln\\,(1+R_h/r_s)-R_h/(r_s+R_h)]$.  For galaxies in halos of $M_h=10^{11}\\,{\\rm M}_\\odot$, which is typical of the mass range probed by our sample, we calculate a launch velocity of $v_{\\rm out}\\approx 170$ \\kms\\ at $r_i=1$ kpc.  For lower-mass galaxies in halos of $M_h=10^{10}\\,{\\rm M}_\\odot$, we find $v_{\\rm out}\\approx 75$ \\kms\\ at $r_i=1$ kpc.  For higher-mass galaxies in halos of $M_h=10^{12}\\,{\\rm M}_\\odot$, we find $v_{\\rm out}\\approx 330$ \\kms\\ at $r_i=1$ kpc.  Naturally, a larger launch speed is required in more massive halos in order to reach out to a larger distance.  Because more massive galaxies on average are forming stars at a higher rate (e.g.\\ Salim \\etal\\ 2007), it is conceivable to have a higher outflow speed from starburst driven winds in a more massive galaxy.  With a larger sample, we will be able to examine whether/how the metal-line boundary depends on star formation activity.  An alternative explanation for the observed finite metal-line boundary at $d\\approx 0.7\\,R_h$ is the critical radius below which cool clouds can develop and stabilize in a multi-phase medium (e.g.\\ Maller \\& Bullock 2004; McCourt \\etal\\ 2012).  Specifically, McCourt \\etal\\ (2012) showed that hot halos can develop extended multiphase structures if the cooling time $t_{\\rm cool}$ is comparable to or less than the dynamical time $t_{\\rm ff}$.  Given that $t_{\\rm cool}\\propto 1/\\rho$ and $t_{\\rm ff}\\propto 1/\\sqrt{\\rho}$, the ratio $t_{\\rm cool}/t_{\\rm ff}$ decreases rapidly with decreasing radius.  While the details depend on the exact density profile of the hot halo, it is conceivable that such condition can be reached at $\\approx 0.7\\,R_h$.  Both the turn-around radius of starburst outflows and cooling radius of a hot halo can provide a physical explanation for the observed metal-line boundary.  Because the former requires on-going star formation to provide the energy input, one can in principle distinguish between the two scenarios by examining whether or not the presence of a metal-line boundary depends on the star formation activity in the galaxies.  \\subsection{The CGM at Low and High Redshifts}   When comparing a luminosity-normalized spatial extent of the CGM at different redshifts, Chen (2012) found that no distinctions can be made between low- and high-redshift galaxies that have dissimilar star formation histories.  The adopted luminosity normalization was meant to account for a likely luminosity-size scaling relation commonly seen in optical disks (e.g.\\ Cameron \\& Driver 2007).  This approach was motivated by the expectation that the CGM is regulated by various complicated physical processes, including gravitational motion and star formation feedback.  Normalizing the observed projected distance by the intrinsic luminosity offers a means of accounting for the intrinsic size differences of individual galactic halos, and represents the first step toward establishing a clear understanding of the processes that drive the observed CGM properties.  The results will facilitate future studies that consider the effect of additional processes such as star formation feedback.  But while rest-frame $B$-band luminosity ($L_B$) is a good proxy of mass as shown in Chen \\etal\\ (2010b), it is an indirect measure of galaxy mass.  Our study presented here utilizes a sample of galaxies with known stellar masses and normalizes the observed extent of halo gas by the inferred dark matter halo radius under a simple spherical overdensity approximation (e.g.\\ Maller \\& Bullock 2004).  The finding in \\S\\ 4.3 is surprisingly very different from the finding of Chen (2012).  Specifically, when accounting for a size-scaling based on halo mass (instead of rest-frame $B$-band luminosity), high-redshift starburst galaxies exhibit both systematically stronger mean absorption strength at a fixed $R_h$-normalized projected distance at $d/R_h\\apl 0.7$ and more extended chemically-enriched gas to $d\\apg R_h$ than low-redshift galaxies.  Furthermore, the differences in the mean absorption strengths of the low- and high-redshift CGM are most pronounced in low-ionization species traced by C\\,II and Si\\,II absorption lines, suggesting distinct ionization conditions between the CGM at $z\\approx 0$ and at high redshifts near starburst galaxies.  We note, however, that a fundamental difference between the low- and high-redshift measurements is the treatment of galaxies with close neighbors.  In our study presented here, we have considered only ``isolated'' galaxies with no known close neighbors spectroscopically identified within 500 kpc in projected distance and $|\\Delta\\,v|<500$ \\kms\\ in velocity separation.  No such distinction was made in the high-redshift sample, and therefore the high-redshift CGM measurements include both isolated galaxies and those with neighbors.  It has been suggested that galaxies in group environments display a more extended, chemically enrichment halo gas than galaxies in the field (e.g.\\ Chen \\etal\\ 2010a; Bordoloi \\etal\\ 2011).  At the same time, the overall cool gas content is also found suppressed in galaxy groups (e.g.\\ Chynoweth \\etal\\ 2011).  It is not clear that the more extended, strong CGM absorption around high-redshift starburst galaxies is due to a signigicant contribution of galaxies in group environments.  A detailed study of the CGM in different group environments at low redshifts will be presented in a separate paper (Liang \\etal\\ 2014 in preparation).  In addition, Steidel \\etal\\ (2010) noted that the Si\\,II absorption lines are saturated at $d\\apl 40$ kpc in the stacked low-resolution spectra.  On the other hand, the ionic transitions observed in individual galaxies at low redshifts are not commonly saturated. Steidel \\etal\\ interpreted the observed declining absorption equivalent width (at least from $d=0$ to $d=40$ kpc) as likely due to a combination of decreasing line-of-sight velocity spread and decreasing covering fraction of the absorbing gas.  In contrast, observations individual galaxies at low redshift indicate that the declining mean absorption equivalent width in stacked spectra is best explained by a declining gas covering fraction along with declining optical depth with increasing projected distance.  Consequently, interpreting the observed difference between low- and hig-redshift observations becomes more complicated.  At $d>40$ kpc, however, Steidel \\etal\\ (2010) also pointed out that the ionic absorption lines appeared to become unsaturated around high-redshift starburst galaxies.  For unsaturated lines in stacked spectra, the observed declining equivalent width is best explained by a declining gas covering fraction along with a declining optical depth, similar to the low-redshift observations.  Comparisons between our means stacked spectra and those from Steidel \\etal\\ could therefore be made for observations at $d>40$ kpc.  Remaining caveats include redshift uncertainties. Individual spectra of these high-redshift galaxies are of low $S/N$ and low spectral resolution with ${\\rm FWHM}\\approx 350$ \\kms.  Random motions between galaxies and the absorbing gas (e.g.\\ Figure 6) can smooth out the signal.  Chen (2012) performed a simple Monte Carlo simulation to show that these uncertainties can result in an overall reduction of the absorption strength by $\\approx 20$\\%, increasing the discrepancy between low- and high-redshift samples.  At the same time, uncertainties in sky subtraction of low $S/N$ spectra may also bias the continuum estimate that is difficult to assess.  It is expected that these uncertainties can be resolved with absorption spectroscopy of background QSOs near individual starburst galaxies (e.g.\\ Rudie \\etal\\ 2012; Turner \\etal\\ 2014).    \\begin{figure*} \\centering \\includegraphics[scale =0.45]{w_civ_rs} \\includegraphics[scale =0.45]{wms_CIV} \\caption{\\civ\\ radial absorption profile normalized by the scale radius $r_s$ of the host dark matter halos ({\\it left}) or by $M_{\\rm star}$ of the absorbing galaxies ({\\it right}).  The scale radius $r_s$ represents a characteristic radius where the dark matter density profile resembles an $r^{-2}$ profile (Navarro \\etal\\ 1997).  We also include in the left panel measurements from stacked spectra using the full galaxy sample (blue pentagons).  In addition, the scaling coefficient of $M_{\\rm star}$ is not a best fit to the \\civ\\ data, but adopted from the best-fit coefficient of Mg\\,II absorbers from Chen \\etal\\ (2010b).  Nonetheless, the result already reveals a clear \\civ\\ enriched halo out to $M_{\\rm star}$-normalized projected distance of $\\approx 160$ kpc (consistent with the finding of $L_B$-normalized \\civ\\ profile from Chen 2012).  In both panels, the \\civ\\ radial absorption profiles at $\\langle\\,z\\,\\rangle=0.041$ (solid points) and at $\\langle\\,z\\,\\rangle=2.2$ (green crosses) appear to be better aligned than what is seen in the bottom right panel of Figure 7 \\& 8. } \\end{figure*}   \\subsection{Uncertainties due to Pseudo-evolution of Halo Radius $R_h$}  We note, however, that a fundamental caveat in comparing $R_h$-normalized CGM profiles at different redshifts is in how $R_h$ is calculated.  Equations (2) and (3) outlines a commonly adopted approach to calculate halo mass $M_h$ and halo radius $R_h$ based on a spherical overdensity model.  A dark matter halo is defined as a collapsed and virialized object with a mean enclosed density equal to $\\Delta_h$ times the mean background density $\\langle\\rho_m\\rangle$ at the redshift of interest.  By this definition, however, as the universe expands, $\\langle\\rho_m\\rangle$ decreases.  Consequently, both $M_h$ and $R_h$ would increase with time even for a static halo. Such pseudo-evolution of dark matter halos has been investigated by several groups, and a consistent finding in numerical simulations is that most halos of $M_h\\apl 10^{12}\\,M_\\odot$ completed matter accretion by $z\\sim 1$ with little/no net growth in mass over the last seven billion years (e.g.\\ Prada \\etal\\ 2006; Diemand \\etal\\ 2007; Diemer \\etal\\ 2013).  Therefore, applying an $R_h$ normalization may impose a pseudo-evolution in the properties of galaxies residing in these low-mass halos.     Given the uncertainties in $R_h$, we explore different scaling relations for an accurate comparison of the CGM absorption profiles in different epochs.  In Figure 13, we present \\civ\\ radial absorption profile normalized by the scaled radius $r_s$ of the host dark matter halos (left panel) or by $M_{\\rm star}$ of the absorbing galaxies (right panel).  The scale radius $r_s$ represents a characteristic radius where the dark matter density profile resembles an $r^{-2}$ profile (Navarro \\etal\\ 1997).  The $r_s$ normalization is motivated by the understanding that when a halo stops accreting, the scaled radius $r_s$ remains constant while the halo mass $M_h$ and halo radius $R_h$ continue to increase as the universe expands and $\\langle\\rho_m\\rangle$ descrases.  As a result, the halo concentration $c\\equiv R_h/r_s$ would also increase.  Indeed, halo concentrations are found to increase rapidly for halos of $M_h\\apl 10^{12}\\,M_\\odot$ at $z\\apl 1$ (e.g.\\ Zhao \\etal\\ 2009).  To calculate $r_s$, we adopt the mean $c$-$M_h$ relation of Zhao et al.\\ (2009), which parameterizes $c$ as a function of halo mass and redshift, $c(M_h, z)$.  The left panel of Figure 13 shows that the $r_s$-normalized \\civ\\ radial profile at $\\langle\\,z\\,\\rangle=2.2$ declines rapidly at $d\\approx (4-6)\\,r_s$.  This is consistent with observations of the $\\langle\\,z\\,\\rangle=0.041$ sample that display a covering fraction of \\civ\\ absorbing gas at $d < 6\\,r_s$ with $\\mkap_{\\rm C\\,IV}\\approx 50$\\% and no detections found beyond $8\\,r_s$.  The agreement in the $r_s$-normalized CGM absorption profiles at low and high redshifts is in stark contrast to the distinction seen in the $R_h$-normalized CGM profiles presented in Figures 7 \\& 8.  We note, however, an inherent uncertainty in $r_s$ as a result of the scatter in the mean $c$-$M_h$ relation ($\\sim\\,0.15$ dex).  Nonetheless, it underscores the fundamental difficulty in making direct comparisons of galaxies and their halo gas properties in two different epochs.  A pseudo-evolution in dark matter halo mass (and therefore halo radius $R_h$) has important implications beyond a constant CGM radial absorption profile.  It would also suggest that chemical enrichment have reached out to the ``edge'' of the dark matter halo ($\\sim 4\\,r_s$) or likely beyond, although a clear border for the observed metal-line absorbers outside of halo radius still implies a finite extent of chemical enrichment around galaxies.  To quantify the extent of chemical enrichment, we will rely on numerical simulations as a guide for determining the physical size of the underlying dark matter halos of our galaxies.  Such study is beyond the scope of this paper. We defer the analysis to a future paper (Liang et al.\\ 2014, in preparation).  Given the uncertainty in $R_h$ and $r_s$ from dark matter halo models, we return to applying empirical properties of the galaxies for scaling halo gas properties.  Instead of $B$-band luminosity, however, we consider normalizing the projected distance of galaxies by stellar mass, $M_{\\rm star}$.  The right panel of Figure 13 shows the $M_{\\rm star}$-normalized \\civ\\ radial profiles for the high- and low-redshift galaxy samples.  We note that the adopted scaling coefficient of $M_{\\rm star}$ is not a best fit to the \\civ\\ data, but adopted from the best-fit coefficient of Mg\\,II absorbers from Chen \\etal\\ (2010b).  Nonetheless, the result already reveals a clear \\civ\\ enriched halo out to $M_{\\rm star}$-normalized projected distance of $\\approx 160$ kpc (in an excellent agreement with the finding of $L_B$-normalized \\civ\\ profile from Chen 2012).  Similar to the $r_s$-normalized profiles in the left panel, the $M_{\\rm star}$-normalized \\civ\\ radial absorption profiles at $\\langle\\,z\\,\\rangle=0.041$ and at $\\langle\\,z\\,\\rangle=2.2$ appear to be well aligned.  With in mind the uncertainty in $R_h$ and the agreement in $M_{\\rm star}$-normalized CGM radial profile, we therefore conclude that no strong evidence has been found to show a vastly different spatial distribution in the CGM absorption properties, despite a completely disparate star formation properties."
    },
    "1402/1402.1788_arXiv.txt": {
        "abstract": "We present the smoothed-particle hydrodynamics implementation SPHGal, which combines some recently proposed improvements in GADGET. This includes a pressure-entropy formulation with a Wendland kernel, a higher order estimate of velocity gradients, a modified artificial viscosity switch with a modified strong limiter, and artificial conduction of thermal energy. With a series of idealized hydrodynamic tests we show that the pressure-entropy formulation is ideal for resolving fluid mixing at contact discontinuities but performs conspicuously worse at strong shocks due to the large entropy discontinuities. Including artificial conduction at shocks greatly improves the results. In simulations of Milky Way like disk galaxies a feedback-induced instability develops if too much artificial viscosity is introduced. Our modified artificial viscosity scheme prevents this instability and shows efficient shock capturing capability. We also investigate the star formation rate and the galactic outflow. The star formation rates vary slightly for different SPH schemes while the mass loading is sensitive to the SPH scheme and significantly reduced in our favored implementation.  We compare the accretion behavior of the hot halo gas. The formation of cold blobs, an artifact of simple SPH implementations, can be eliminated efficiently with proper fluid mixing, either by conduction and/or by using a pressure-entropy formulation. ",
        "introduction": "Smoothed-particle hydrodynamics (SPH) is a numerical method for solving fluid equations in a Lagrangian fashion \\citep{1977AJ.....82.1013L, 1977MNRAS.181..375G}. It has found a widespread use in a variety of astrophysical problems (e.g. \\citealp{2010ARA&A..48..391S,2012JCoPh.231..759P}). The success of SPH lies in the exact conservation of physical properties and the adaptive resolution that traces mass. In addition, the Lagrangian nature ensures Galilean invariance by construction. Practically, the method is also straightforward to be incorporated with tree-based gravity solvers and is intuitive for including sub-resolution physics.  However, recent studies (e.g. \\citealp{2007MNRAS.380..963A, 2010MNRAS.401..791S}) have shown that standard SPH has serious difficulties to properly model fluid mixing. Its accuracy is therefore compromised when modeling the interstellar medium \\citep{2011MNRAS.415..271H} and galaxy formation \\citep{2012MNRAS.424.2999S, 2013MNRAS.429.3353N}. The problem originates from a numerical artifact at contact discontinuities (often referred to as the spurious 'surface tension'). One solution proposed by \\citet{2008JCoPh.22710040P} is to include artificial conduction to alleviate such artifacts, though with the potential issue of being overly diffusive. Therefore, several conduction switches have been suggested to reduce unwanted conduction away from the entropy discontinuities \\citep{2012MNRAS.422.3037R, 2012A&A...546A..45V}. An alternative solution to the problem is to modify the definition of density so that the pressure is smoothed by construction \\citep{2001MNRAS.323..743R, 2010MNRAS.405.1513R, 2013ApJ...768...44S, 2013MNRAS.428.2840H}. To its advantage no extra dissipation is needed, the entropy is still conserved and the contact discontinuities remain sharp.  Another criticism of SPH is its slow convergence rate. SPH is second order accurate only in the continuous limit. When the fluid is discretized into SPH particles, there exists a zeroth order error in the equation of motion \\citep{2010MNRAS.405.1513R, 2012JCoPh.231..759P}. This error can be reduced by using more neighboring particles within the kernel.  However, the commonly used cubic spline kernel is subject to the pairing instability when too many neighboring particles are used \\citep{2012JCoPh.231..759P}. \\citet{2010MNRAS.405.1513R} introduced a new kernel function without an inflection point to prevent the pairing instability. \\citet{2012MNRAS.425.1068D} proposed to use the Wendland $C^4$ kernel and demonstrated its stability despite having an inflection point.  These studies demonstrate a significantly improved convergence rate compared to the standard SPH with only a minor impact on the computational efficiency.  With these recent developments, many authors have updated the existing 'problematic' implementations to improve the accuracy and ensure the scientific credibility of their simulations. \\citet{2013arXiv1307.0668P} presented simulations of the formation of an idealized (non-radiative) galaxy cluster using three different methods: the traditional SPH, their improved implementation (SPHS) and an adaptive mesh refinement (AMR) code (see also \\citet{2013MNRAS.432..711H} for a recent AMR/SPH comparison). They find that SPHS and AMR are in excellent agreement while the traditional SPH shows a different behavior. On the other hand, \\citet{2013arXiv1311.2073H} presented cosmological simulations including more complicated physical processes and found very little difference between their improved SPH implementation and the traditional SPH. \\citet{2013arXiv1312.1391F} simulated the accretion of supermassive black holes at high redshifts and found that their improved SPH gives rise to slightly faster black hole growth, while the star formation is unaffected.  In this paper, we incorporate the most important improvements into the hydrodynamic code GADGET-3 \\citep{2005MNRAS.364.1105S} and test different implementations with idealized hydrodynamic tests. We use the pressure-entropy formulation presented in \\citet{2013MNRAS.428.2840H} and the Wendland $C^4$ kernel with 200 neighboring particles.  The artificial viscosity scheme we adopt resembles \\citet{2010MNRAS.408..669C} but with a slight modification in its functional form. We also include artificial conduction similar to \\citet{2012MNRAS.422.3037R}. However, the major role of conduction in our case is to help capture shocks. We show that the pressure-entropy formulation performs poorly in the Sedov explosion test without artificial conduction. At contact discontinuities, however, the pressure-entropy formulation alone is able to properly model the fluid mixing. We therefore suppress the artificial conduction in shear flows.   We also investigate the properties of different SPH schemes in a more complex three dimensional simulation of isolated disk galaxy models with gas, star formation and supernova feedback. Standard SPH implementations can result in the formation of artificial large, kpc-sized, holes which are then sheared apart. If too much viscosity is introduced the gaseous disks become unstable and the simulations fail dramatically.  Fortunately, such instabilities can be avoided with a slight modification of the viscosity limiter.  We present a modified SPH version, which we term SPHGal, that passes the Gresho, Sod shock tube, Sedov explosion, 'square', Keplerian ring, Kelvin-Helmholtz and 'blob' test and also performs well in more realistic galaxy simulations. The strengths and limitations of this implementation are discussed in detail.    This paper is organized as follows: in Section \\ref{sec:method} we present the details of our SPH implementation. The results of the idealized hydrodynamic tests are shown in Section \\ref{sec:hydrotest} with detailed discussions of the pros and cons of different implementations. In Section \\ref{sec:disk} we present simulations of isolated disk galaxy models and show the properties of the gaseous disk, the star formation rate and galactic outflow, and the accretion behaviors. We summarize and discuss our work in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper we study the performance of different algorithmic implementations of the SPH code GADGET for a variety of idealialized hydrodynamic tests including the Gresho, Sod shock tube, Sedov explosion, 'equilibrium square', Keplerian ring, Kelvin-Helmholtz, and the 'blob' test (Section \\ref{sec:hydrotest}). We also test the impact on the dynamical evolution of a Milky Way like disk galaxy - including cooling from a hot gaseous halo - as well as a gas-rich high-redshift disk galaxy including metal cooling, metal enrichment, star formation and feedback from supernova explosions.  We test a density-entropy and a pressure-entropy formulation in combination with recently proposed implementations for the treatment of artificial viscosity and conduction (Section \\ref{sec:method}). The study indicates that implementations with a pressure-entropy formulation \\citep{2013ApJ...768...44S,2013MNRAS.428.2840H} in combination with a Wendland $C^4$ kernel with 200 neighbors, variable artificial viscosity with higher order shock detection \\citep{2010MNRAS.408..669C} and artificial conduction \\citep{2008JCoPh.22710040P, 2012MNRAS.422.3037R} pass all idealized hydrodynamic tests.  We propose a modified implementation for artificial viscosity with a strong viscosity coefficient limiter with high sensitivity to the velocity curl (Section \\ref{sec:AV}). In addition to passing all the idealized tests, our fiducial implementation performs very well for the disk evolution simulations, where some other implementations can develop a viscous instability (Section \\ref{sec:disk}). Despite the successes our fiducial code still shows weaknesses in low Mach number shear flows and the rate of convergence. Still, the above results indicate improvements with respect to most standard SPH implementations (e.g. \\citealp{2005MNRAS.364.1105S,2009ApJS..184..298W,2011A&A...529A..27H}). \\textbf{ These improvements, along with the advantages discussed in the introduction, keep SPH a competitive numerical method. Comparison between results of SPH and grid-based methods (e.g. \\citealp{2002A&A...385..337T}; \\citealp{2014ApJS..211...19B}) can also provide important information as they have very different numerical artifacts and are highly complementary. }     The actual algorithmic implementation has a significant impact on the gas outflow properties of our disk evolution models. Implementations with better fluid mixing capabilities in general result in reduced mass loading. In the Milky Way case, our fiducial implementation has about an order of magnitude lower mass loading compared to a standard implementation, while at the same time the star formation rate is only slightly higher. For the gas-rich, high-redshift disk model, the mass loading becomes less sensitive (but still notable) to the adopted SPH scheme. In addition, we demonstrate that either artificial conduction or a pressure-entropy formulation (or both) suppresses the formation of artificial cold blobs. \\textbf{ We note that the star formation model as well as the feedback recipe have also significant impact on the evolution of galaxies (e.g., \\citealp{2010MNRAS.409.1541S}, \\citealp{2012MNRAS.423.1726S}). However, an accurate scheme for the hydrodynamics is equally (if not more) important, as the star formation and feedback models need to be calibrated under the 'correct' hydrodynamics. The impact of star formation and feedback models will be addressed in future work. }   In Table \\ref{table:sum}, we present a summary of all tests in this paper. SPH implementations marked with '$\\checkmark$' passed the respective test while those marked with '$\\times$' failed. Again, these tests indicate that SPH can perform well with a more accurate kernel, a pressure-entropy formulation, artificial viscosity with a strong, curl sensitive limiter and a higher order shock indicator and artificial conduction. Our fiducial implementation, which we term SPHGal, passes all tests. \\footnote{ We find (not shown) that the weak limiter scheme 'pe-avwl-ac' also passes all of the idealized tests that might show differences with the strong limiter (i.e. shock tube, Sedov, Keplerian ring and the blob test). It only fails in the disk simulation, where both shocks and shear flows coexist.}      \\begin{table} \\caption{Summary of the results of the idealized hydrodynamic tests in Section \\ref{sec:hydrotest} and the galactic disk test in Section \\ref{sec:disk}. SPH implementations marked with $\\checkmark$ passed the test while those marked with $\\times$ failed the test. Our fiducial implementation 'pe-avsl-ac' passed all tests.} \\label{table:sum} \\begin{tabular}{| l | l | l |} \\hline\\hline Test&                 SPH scheme&                 Results      \\\\ \\hline Gresho&               de-avB (cs64)&                $\\times$ \\\\%(slow convergence rate)   \\\\ &               de-avsl-ac (wld200)&          $\\checkmark$   \\\\ &               pe-avsl-ac (wld200)&          $\\checkmark$    \\\\ \\hline Shock tube&           pe-avsl-ac&                   $\\checkmark$     \\\\ &           pe-avwl-ac&                   $\\checkmark$     \\\\ \\hline Sedov     &           de-avsl&                      $\\checkmark$     \\\\ &           pe-avsl&                      $\\times$  \\\\%(large scatter)   \\\\ &           de-avsl-ac    &               $\\checkmark$     \\\\ &           pe-avsl-ac    &               $\\checkmark$     \\\\ &           pe-avwl-ac    &               $\\checkmark$     \\\\ &           pe-avsl-ac-erho&              $\\times$  \\\\%(shock travels too slowly)  \\\\ \\hline Keplerian ring &      de-avB-lvg       &            $\\times$  \\\\%(noisy velocity gradient)   \\\\ &      pe-avsl-ac-lvg   &            $\\times$  \\\\%(noisy velocity gradient)   \\\\ &      de-avB           &            $\\times$  \\\\%(too much viscosity)  \\\\ &      pe-avsl-ac       &            $\\checkmark$    \\\\ &      pe-avwl-ac       &            $\\checkmark$    \\\\ \\hline Hydrostatic    &      de-avsl-ac       &            $\\times$  \\\\%(deformed by surface tension) \\\\ &      pe-avsl          &            $\\checkmark$     \\\\ &      pe-avsl-ac       &            $\\checkmark$     \\\\ \\hline Kelvin-Helmholtz  &   pe-avsl-ac       &            $\\checkmark$     \\\\ \\hline Blob       &   pe-avsl-ac       &            $\\checkmark$     \\\\ &   pe-avwl-ac       &            $\\checkmark$     \\\\ \\hline Galactic disk &           de-avB&                      $\\times$     \\\\ &           de-avwl&                      $\\times$  \\\\ &           de-avwl-ac    &               $\\times$     \\\\ &           de-avsl    &               $\\checkmark$     \\\\ &           de-avsl-ac &              $\\checkmark$  \\\\ &           pe-avB&                      $\\checkmark$     \\\\ &           pe-avwl&                      $\\times$  \\\\ &           pe-avwl-ac    &               $\\times$     \\\\ &           pe-avsl    &               $\\checkmark$     \\\\ &           pe-avsl-ac &              $\\checkmark$  \\\\ \\end{tabular} \\end{table}  In the follwing we discuss our main results in more detail:  \\begin{itemize} \\item With the PE formulation the spurious surface tension at contact discontinuities is eliminated by construction and fluid mixing can be modeled properly without the help of AC \\citep{2010MNRAS.405.1513R,2013ApJ...768...44S, 2013MNRAS.428.2840H}. We therefore suppress AC in shear flows to prevent over-mixing by including a quadratic limiter similar to the Balsara switch (Section \\ref{sec:ac}).  Another advantage of this limiter is that it avoids unwanted conduction in a self-gravitating system in hydrostatic equilibrium.   \\item It is necessary to include the AC in the PE formulation if strong shocks are involved (Section \\ref{sec:sedov}). The weakness of the PE formulation stems from the entropy-weighted sum in the pressure estimate, which makes the results noisy and even biased if the entropy variation is large from one particle to another. Including AC greatly improves the results in the Sedov explosion test. In addition, the entropy-weighted density gives a biased estimate (the pre-shock 'bump') in the Sedov explosion test even when the particles are distributed regularly. The mass-weighted density is a more reliable estimate and should be used whenever the information of density is  needed. Indeed, the only entropy-weighted quantity we use is the pressure estimate in the equation of motion. Using the entropy-weighted density in the dissipation terms leads to an incorrect prediction of the shock position.    \\item We find a feedback-driven instability developing in a typical isolated disk galaxy if the gaseous disk becomes too viscous.  The hot bubbles created by SN explosions are too viscous to be sheared away by the differential rotation.  They merge with one another and eventually form large holes in the disk.  The AV switch with a strong limiter is able to suppress the viscosity so that the disk remains stable.  On the other hand, it could also be possible that it suppresses viscosity too much and cannot capture shocks properly.  However, the Sedov explosion test is recovered equally well and provides credibility for our modified AV scheme.  In the cases where both shocks and shear flows coexist, it is in fact difficult to determine how much viscosity is appropriate. Compromises have to be made between properly modeling shocks and avoiding artificial shear viscosity. This seems to be a general issue for most SPH schemes that use AV for shock capturing (an exception might be 'Godunov SPH', e.g. \\citealp{2010MNRAS.403.1165C}, \\citealp{2011MNRAS.417..136M}). Adopting the PE formulation and including AC alleviates the situation, which suggests that spurious surface tension also plays a role for the instability. The boundaries between the hot bubbles and the cold ISM are sustained by the spurious surface tension in the DE formulation. As such, they are difficult to destroy once being created.  \\item We have investigated the star formation rate as well as the mass loading of the disk galaxy models with different SPH schemes. The SFR increases slightly when we adopt the SPH schemes that allow more efficient fluid mixing. Here the hot gas in the disk (created by SN explosions) is continuously mixed with the large amount of cold gas and therefore cools faster.  For this reason, the amount of cold gas available to star formation is higher. The mass loading is affected even more and drops by almost an order of magnitude (in the MW case) from a standard SPH implementation to our fiducial implementation. This is partly related to the increased SFR that transforms some of the gas into stars.  In addition, the continuous mixing with the ambient ISM weakens the driving force of the shock-heated particles and hence decreases the mass loading. This is in contrast to \\citet{2013arXiv1311.2073H} who have made similar improvements to their SPH scheme and claim that they in general find little difference in SFR as well as the mass loading using different SPH schemes. \\textbf{ The discrepancy may arise from the different feedback models. It is expected that the thermal feedback would be more sensitive to the fluid mixing, as mixing provides an extra channel for the hot gas to cool and thus weakens the feedback. However, we note that the kinetic feedback would still convert some fraction of kinetic energy into thermal energy via shock heating and therefore is not completely unaffetced by mixing. } In the case of the gas-rich, high-redshift disk model, the mass loading becomes less sensitive to the adopted SPH scheme, which is in better agreement with \\citet{2013arXiv1311.2073H}. One possible explanation is that the higher star formation rate drives the outflow more efficiently despite the increased capability of fluid mixing and hence weakening the differences.     \\item We investigate the accretion behavior of the hot gas halo onto the Milky Way like disk. As commonly being criticized in the literature, the traditional SPH (DE without conduction) generates plenty of small blobs due to the lack of proper fluid mixing. Including AC efficiently eliminates these blobs, in agreement with \\citet{2013MNRAS.434.1849H}. The PE formulation, with or without AC, is also capable of eliminating the blobs. In addition, depending on the presence of AC, the accretion morphology shows notable difference: inflowing gas forms smooth filamentary structures if AC is included. This might originate from the different mixing mechanism between AC and PE formulation. In PE formulation, mixing is a dissipation-less, entropy-conserving process that relies on turbulent motion. Hot and cold gas do not mix if the turbulent motion is weak. On the other hand, AC promotes mixing by explicitly smoothing out entropy differences, leading to smoother morphology.      \\end{itemize}  In summary, we have presented an updated SPH implementation SPHGal that performs more accurately in several idealized hydrodynamic tests. The outstanding problem in the traditional SPH, i.e. the poor capability of fluid mixing, no longer exists thanks to the PE formulation.  In addition, we have a controllable diffusion mechanism (the artificial conduction) similar to the implicit numerical diffusion in the grid-based methods. The higher order velocity gradients prevent false triggering of AV. The AV switch, which involves a strong limiter, is able to capture shocks properly and also to avoid too much viscosity in a disk galaxy. The convergence rate is significantly improved (still slower than the grid-based methods) and shows no sign of saturation with increasing resolution, though the intrinsic issue remains when we move further into the subsonic regimes. We conclude that, with all the modifications, SPHGal is an accurate and valuable numerical method for galaxy formation simulations and many other astrophysical applications."
    },
    "1402/1402.3744_arXiv.txt": {
        "abstract": "{ We use a type 2 AGN sample from SDSS DR7 in which the \\Oiii\\ $\\lambda$5007 emission line can be modeled by two Gaussian components, a broad wing plus a narrow core, to investigate the origin of the broad wing and the connection between the velocity shift of the broad wing and the physical parameters of active galactic nuclei (AGNs) as well as their host galaxies. We find that the flux of the wing components is roughly equal to that of the core components in statistic. However, the velocity shift of the wing component has only weak, if any, correlations with the physical properties of AGNs and the host galaxies such as bolometric luminosity, the Eddington ratio, the mass of supermassive black holes, D4000, $\\rm H\\delta_A$ or stellar mass. Comparing the velocity shift from our type 2 AGN sample to that from type 1 sample in \\citet{zhang2011} , we suggest the \\Oiii\\ broad wing originates from outflow. ",
        "introduction": "Active galactic nuclei (AGNs) feedback now appears as a crucial process in galaxy formation and evolution. It is well-known that the tight correlation between nuclear black hole mass ($M_{\\rm BH}$) and bulge stellar velocity dispersion (i.e., the $M_{\\rm BH} - \\sigma_*$ relation; \\citealt{ferrarese2000}; \\citealt{gebhardt2000}; \\citealt{tremaine2002}) is compelling evidence for a close connection between the evolution of supermassive black holes and their host galaxies. AGN feedback is the most likely explanation for this relation (e.g., \\citealt{silk1998}; \\citealt{fabian1999}). \\citet{hopkins2006} also indicate the importance of feedback in the evolutionary model for starbursts, AGN activity and spheroidal galaxies. This feedback can terminate star formation in the host galaxy and cease gas accretion onto the nuclear black hole, with the form of radiation, winds, jets and outflows.  The various emission lines from narrow-line regions (NLRs) of AGN are ideally suited to study the central regions of AGNs, as well as the interaction between the central engine and its host galaxy. Unlike the broad-line regions (BLRs), the NLR is spatially resolvable, at least for nearby galaxies. It is generally believed that the narrow emission lines are produced by clouds illuminated by the central AGN, and the kinematics of the NLR clouds are mainly dominated by the gravitational potential of the bulge (e.g.,  \\citealt{whittle1992}; \\citealt{nelson1996}). Since the NLR connects to various factors such as the energy input from the central engine, the structure of AGN, radio jets, and star formation etc, it accesses a number of key questions of AGN phenomenon.  \\Oiii\\ $\\lambda$$\\lambda$4959,5007 emission lines are commonly used to study the properties of NLRs. It is usually the strongest narrow line in AGNs at optical band and cleanly isolated from other emission and absorption features in the optical spectrum. The line profile of \\Oiii\\ doublets in low-redshift AGNs is usually asymmetric. In most case, there is a sharper fall-off to the red than to the blue, and the redshift of \\Oiii\\ is negative compared to the system velocity derived from different indicators, such as the low-ionization lines (\\Nii, \\Sii), or stellar absorption lines (\\citealt{heckman1981} and consequent researches). This asymmetry feature of \\Oiii\\ line has been suggested as an indicator of outflows in Seyfert 1s and 2s (\\citealt{heckman1981}; \\citealt{whittle1988}; \\citealt{colbert1996}; \\citealt{crenshaw2010}), Narrow Line Seyfert 1 galaxies (\\citealt{bian2005}; \\citealt{komossa2008}), type 1 quasars (\\citealt{heckman1984}; \\citealt{boroson2005}) and narrow line radio galaxies (\\citealt{holt2008}).  It is now believed that a broad blue wing in addition to the main narrow component is ubiquitous in \\Oiii\\ emission. Many previous studies suggest this blue wing is contributed by AGN outflows, they also attempt to correlate the parameters of lines, such as the \\Oiii\\ blueshift or/and the \\Oiii\\ line width, with the physical properties of AGNs (\\citealt{zamanov2002}; \\citealt{aoki2005}; \\citealt{boroson2005}; \\citealt{bian2005}; \\citealt{komossa2008}). Based on homogeneous samples of radio-quiet Seyfert 1 galaxies and QSOs selected from SDSS, \\citet{zhang2011} find that the blueshift of \\Oiii\\ has only weak correlations with fundamental AGN parameters, such as the nuclear continuum luminosity at 5100\\AA\\ ($L_{5100}$), black hole mass ($M_{\\rm BH}$), and the Eddington ratio ($L_{\\rm bol}/L_{\\rm Edd}$). \\citet{alexander2012} mentioned that in statistics, the width and luminosity of the blue wing increase with \\Oiii\\ luminosity, but independent of radio loudness, indicating the outflows are driven by AGN radiation rather than relativistic jets. \\citet{zhang2008} found that Seyfert 1s have lower \\Nii/\\Ha\\ ratios than Seyfert 2s and the location of Seyfert 1s on the BPT diagram varies with extinction of broad lines, suggesting that the inner dense NLR is obscured by dusty torus. The inner dense NLR might be the place where the \\Oiii\\ blue wing originates (\\citealt{zhang2013}). \\citet{stern2013} revisit the location of type 1 and type 2 AGNs on the BPT diagrams, finding similar result as \\citet{zhang2008}---type 1 AGNs are offset to lower \\Sii/\\Ha\\ and \\Nii/\\Ha\\ ratios. However, they conclude that this offset between type 1 and type 2 AGNs is a selection effect rather than dust extinction.  In this paper, we will explore the asymmetric behavior of \\Oiii~$\\lambda$$\\lambda$4959,5007 lines in more detail, studying the origin of \\Oiii\\ asymmetry. In Section 2, we describe the sample selection and data analysis. We show the results in Section 3. In Section 4, the origin of the broad wing is discussed. Our conclusions are given in Section 5. Throughout this paper, a cosmology with {\\slshape H}$_0$ = 70 \\kms\\ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_\\Lambda = 0.7$ is adopted. ",
        "conclusions": "We select a type 2 AGN sample from SDSS DR7. In this sample, two Gaussian components are required to model the \\Oiii\\ $\\lambda$5007 emission line, a broad wing plus a narrow core. We measure the velocity shift (relative to the absorption lines), line width and flux of both components. Combining our type 2 AGN sample with a type 1 sample from \\citet{zhang2011}, we find that: \\begin{enumerate} \\item {there is a tight correlation between the fluxes of wing and core components in  both type 1 and type 2 samples. In both samples, the flux of the wing components is roughly equal to that of the core components.} \\item {in the unification scheme of AGNs, the type 1 and 2 AGNs are intrinsically the same, their different appearance is due to that we observe them from different direction, a dusty torus blocks the continuum source and broad-line region in type 2 AGNs. The difference in the velocity shift of the broad wing between type 1 and type 2 AGNs consists with a picture in which the broad wing originates from outflows and the outflows blow out in a direction perpendicular to the accretion disk with a certain opening angle.} \\item {the velocity shift of the wing component has only weak, if any, correlations with the physical properties of AGNs (bolometric luminosity, the Eddington ratio and the mass of supermassive black holes) and the host galaxies (D4000, $\\rm H\\delta_A$ or stellar mass). We suggest the lack of correlation is due to that the outflow is driven by both AGN and star formation activity. However, the contribution of AGN and star formation activity to the outflow varies from object to object. Future IFU survey like MaNGA (Mapping Nearby Galaxies at APO) will help us to understand, AGN and star formation, which is the primarily driver of outflow in a certain object.} \\end{enumerate}  \\normalem"
    },
    "1402/1402.1577_arXiv.txt": {
        "abstract": "Seyfert galaxies have traditionally been classified as radio-quiet active galactic nuclei. A proper consideration of the nuclear optical emission however proves that a majority of Seyferts are radio-loud. Kpc-scale radio lobes/bubbles are in fact revealed in sensitive observations at low radio frequencies of several Seyferts. Through the use of very long baseline interferometry, we have been able to determine the direction of the parsec-scale jets in some of these Seyfert galaxies. The misalignment between the parsec-scale jets and the kpc-scale lobes that is typically observed, is either suggestive of no connection between the two, or the presence of curved jets that power the radio lobes. In this context, we briefly discuss our new low radio frequency GMRT observations of two Seyfert galaxies with lobes. ",
        "introduction": "\\vspace{-0.3cm} Active galactic nuclei (AGN) are the centres of a small fraction of galaxies that harbour actively accreting supermassive black holes (SMBHs; $10^7-10^9~M_\\sun$). Only about $15\\%-20\\%$ of these have relativistic outflows that are ejected from the BH-accretion disk systems which propagate to 10s$-$100s of kiloparsecs, way beyond the extents of their host galaxies \\citep{Kellermann89}. These AGN are referred to as radio-loud. The vast majority of AGN on the other hand, lack kpc-scale collimated radio outflows, making them radio-quiet. The formal quantitative definition for radio-loudness provided by \\citet{Kellermann89} is that the ratio ($R$) of the radio flux density at 5 GHz must exceed the optical $B$-band flux density at 4400$\\AA$~by~$\\ge10$. Seyfert galaxies have historically been identified as radio-quiet AGN (i.e., they have $R<10$). Since radio-loud AGN are typically hosted by large elliptical galaxies, and Seyferts mostly reside in spiral or lenticular  galaxies, explanations for the kpc-scale radio emission have focused on differences in the host galaxy types: for instance, differences in black hole masses and spins, both of which are different in spiral and elliptical galaxies, have been suggested to explain the differences in radio properties \\citep[e.g.,][]{Sikora07}. The inherent assumption in the estimation of $R$ however, is that the $B$-band luminosity is dominated by the AGN. This is unlikely to be true for Seyfert galaxies. \\citet{HoPeng01} extracted the optical nuclear luminosities for a large number of Seyfert galaxies through a proper modelling of their galactic bulge emission in high resolution {\\it HST} images. Remarkably, after taking into account their optical nuclear luminosities, Ho \\& Peng found that a majority of the Seyfert galaxies were radio-loud. \\citet{Kharb14} estimate that $\\ge70\\%$ of Seyfert galaxies belonging to the Extended 12$\\mu$m Seyfert sample \\citep{Rush93} are radio-loud once their optical nuclear emission is considered. Furthermore, sensitive radio observations at low radio frequencies have revealed the presence of radio structures of typical extents 1$-$10 kpc in several Seyfert galaxies. \\citet{Gallimore06} found that $>44\\%$ of Seyferts belonging to the {\\it complete} CfA+12$\\mu$m sample exhibited kpc-scale radio structures when observed with the Very Large Array in the (sensitive) D-array configuration. These structures typically appear like edge-brightened lobes. However, they differ from the lobes of powerful Fanaroff-Riley II (FRII) sources in that they usually do not show collimated jets leading into them or hotspots at their leading edges. Seyfert lobes raise interesting questions. Are they powered by AGN jets? If yes, do very long baseline interferometry (VLBI) observations reveal parsec-scale radio jets in Seyferts, similar to what is observed in radio-loud AGN ? Are these parsec-scale jets aligned (or not) with the kpc-scale radio lobes? These are some of the questions that we have been attempting to address in our study of Seyfert galaxies with lobes. ",
        "conclusions": "\\vspace{-0.3cm} A proper consideration of the nuclear optical emission in Seyfert galaxies reveals that the majority of them are radio-loud. Sensitive radio observations of Seyferts reveal the presence of kpc-scale lobes in them. VLBI observations detect parsec-scale jets in Seyferts that are misaligned by large degrees ($50^\\circ-60^\\circ$) to the kpc-scale lobes. This could either indicate that the jets are not associated with the lobes or that they are curved and indeed power the lobes. The former scenario is disfavoured by several observational findings that indicate that the lobes are AGN outflow rather than starburst wind driven. Curved radio jets could arise due to precession of the central engine or host galaxy ISM rotation and jet-medium interaction. We are currently in the process of analysing 325 and 610~MHz data from the GMRT of two Seyfert galaxies with lobes, {\\it viz.,}, NGC4235 and NGC4594. The spectral ageing analysis that we plan to carry out in the near future will reveal the age of the lobes in these Seyferts, while an analysis of archival VLBI data will reveal the structure of their parsec-scale jets.  \\begin{figure} \\includegraphics[trim = 0mm 12mm 0mm 0mm, width=5.5cm]{fig2.ps} \\hspace*{-0.07cm}\\includegraphics[trim = 5mm 12mm 0mm 0mm, width=10.3cm]{fig3.ps} \\hspace*{2.2cm} \\includegraphics[trim = 0mm 12mm 0mm 0mm, width=10.3cm]{fig1new.ps} \\caption{Top panels: Preliminary 610~MHz contour images of NGC4594 (the Sombrero galaxy) observed with the GMRT. Different contour levels are chosen to highlight the Seyfert lobes (top left) and the lobe+galactic radio emission (top right). Bottom panel: 610~MHz radio contours in white superimposed on the DSS optical image of the host galaxy (Kharb et al., 2014, in prep.).} \\end{figure}  \\vspace{-0.3cm}"
    },
    "1402/1402.6950_arXiv.txt": {
        "abstract": "The MOdified Newtonian Dynamics (MOND) is an attempt to modify the gravitation theory to solve the Dark Matter problem. This phenomenology is very successful at the galactic level. The main effect produced by MOND in the Solar System is called the External Field Effect parametrized by the parameter $Q_2$. We have used 9 years of Cassini range and Doppler measurements to constrain $Q_2$. Our estimate of this parameter based on Cassini data is given by $Q_2=(3 \\pm 3)\\times 10^{-27} \\ \\rm{s^{-2}}$ which shows no deviation from General Relativity and excludes a large part of the relativistic MOND theories. This limit can also be interpreted as a limit on a external  tidal potential acting on the Solar System coming from the internal mass of our galaxy (including Dark Matter) or from a new hypothetical body. ",
        "introduction": "One of the mysteries of modern astrophysics is known as the Dark Matter problem. This problem comes from the discrepancies between observations of galactic systems and the predictions of General Relativity (GR) and the Standard Model of particles. Some galactic and extragalactic dynamical observations can not be explained by GR and by the amount of observed matter.  Three possibilities can be found in the literature to solve this problem: the presence of unseen mass-energy ; a modification to the theory of gravity ; or both~\\cite{famaey:2012fk}.  Modified Newtonian Dynamics (MOND) is an attempt to modify the gravitation theory to solve the Dark Matter problem. This phenomenology proposed by Milgrom 30 years ago~\\cite{milgrom:1983uq,*milgrom:1983kx,*milgrom:1983fk} is very successful in explaining galactic observations (see for examples \\cite{famaey:2012fk,milgrom:2008uq,sanders:2002fk} and references given in the MOND webpage~\\footnote{\\url{http://www.astro.umd.edu/~ssm/mond/litsub.html}}). In particular, MOND superbly explains the galactic rotation curves \\cite{sanders:2002fk} and automatically recovers the Tully-Fisher law which establishes a relation between the luminosity and the rotation velocity of a spiral galaxy \\cite{tully:1977kx,*mcgaugh:2000vn}.  The main idea of MOND theory is to modify the standard Newtonian gravitation law $\\mathbf a=\\mathbf g_N$ (where $\\mathbf a$ is the acceleration of a test particle and $\\mathbf g_N$ is the Newtonian gravitational field) by the relation $\\mathbf a=\\mathbf g$ with $\\mathbf g$ determined by the relation \\begin{equation}\\label{eq:mond_g} \\mu\\left(\\frac{g}{a_0}\\right)\\mathbf g=\\mathbf g_N. \\end{equation} Here, $\\mu$ is a function of the ratio $g/a_0$ between the norm of the gravitational field $g$ and the MOND acceleration scale $a_0$. The usual value of the MOND acceleration scale $a_0$ is between $0.9 \\times 10^{-10}\\ \\rm{m/s^2}$  and $1.2\\times 10^{-10} \\ \\rm{m/s^2}$ \\cite{gentile:2011uq}.  In order to recover the Newtonian dynamics in a strong gravitational field (formally when $g>>a_0$), the function $\\mu$ needs to satisfy $\\mu(x)\\to 1$ when $x\\to\\infty$. The MOND regime appears in the limit of weak gravitational field $g<<a_0$ where the interpolating function has to satisfy $\\mu(x)\\to x$ when $x\\to 0$ in order to explain galactic rotation curves~\\cite{milgrom:1983uq,milgrom:1983kx,milgrom:1983fk}. Various $\\mu$ functions are possible interpolating between the MOND and the Newtonian regimes.  Relativistic extensions of MOND have been developed and are more satisfactory from a theoretical point of view. The Bekenstein Tensor-Vector-Scalar (TeVeS) theory was the first relativistic extension of MOND~\\cite{sanders:1997kx,*bekenstein:2004fk,*sanders:2005fk}.  These relativistic extensions have now evolved into Einstein-Aether theories~\\cite{jacobson:2001qf,*zlosnik:2006dq,*zlosnik:2007bh,*zhao:2007vn}. Very detailed reviews of relativistic MOND theories can be found in~\\cite{bruneton:2007vn,famaey:2012fk}. Finally, a new interpretation of MOND in term of dipolar dark matter has been developed in~\\cite{blanchet:2007kx,*blanchet:2008fv,*blanchet:2008vn,*blanchet:2009kx}.  While being very successful at the galactic level, the modification of the gravitation theory provoked by MOND has to be small enough in the Solar System to be able to recover the well known Solar System dynamics. Within the Solar System, three types of effects can arise due to MOND, the most important one being the MOND External Field Effect\\footnote{We use the term External Field Effect as in \\cite{blanchet:2011ys} to denote this effect that should not be confused with the fact that large external gravitational field may influence the internal dynamics directly because it is the total acceleration that enters the $\\mu$ function \\cite{milgrom:1983fk}} \\cite{milgrom:2009vn,blanchet:2011ys,blanchet:2011zr}.  The EFE produces a quadrupole correction (parametrized by $Q_2$) to the Newtonian potential which increases with the distance to the Sun. As shown in \\cite{blanchet:2011ys,blanchet:2011zr}, Saturn is the Solar System body that is most likely to allow for detection of the EFE, since the orbit is relatively far from the Sun but with orbit period short enough to allow separation of initial conditions from dynamical model parameters. In order to constrain MOND EFE, we consider below radio tracking measurements of the Cassini spacecraft, taking care to account for systematic correlations between measurements and the orbits of the spacecraft and the orbit of Saturn. We have added the MOND EFE potential to the dynamical model of the orbits of the planets and have estimated the $Q_2$ parameter in the global fit. With this approach, we derive an estimate of the parameter characterizing the MOND EFE that constrains very severely the MOND theory. Moreover, the quadrupolar potential parametrized by $Q_2$ can also be produced by a tidal interaction. Therefore, our constraint on $Q_2$ can also be interpreted as a constraint on the internal mass (including Dark Matter) of our galaxy \\cite{braginsky:1992uq,klioner:1993fk}  or as a constraint on the mass of a hypothetical new body \\cite{iorio:2012rm,*iorio:2010yg}. ",
        "conclusions": "In this paper, we used the Earth-Cassini range data in order to estimate the MOND external field effect parameter $Q_2$ that appears in MOND theory based on modified Poisson equation. For this, we enhanced the dynamical model used for data analysis~\\cite{standish:2012fk} to include the effects predicted by MOND given by~(\\ref{eq:efe}).  The resulting estimate of $Q_2$ is given by~(\\ref{eq:Q2est}). Cross validation tests show that the estimated uncertainty in this parameter is sound. No significant deviation from GR favoring MOND  has been detected in this set of data. Moreover, the very stringent constraint (\\ref{eq:Q2est}) excludes the major part of the interval~(\\ref{eq:Q2}) on $Q_2$ computed theoretically in \\cite{blanchet:2011ys}. The constraint on $Q_2$ can be written as a constraint on $q$ (\\ref{eq:qest1}-\\ref{eq:qest2}) which can directly be compared to the Table I of \\cite{milgrom:2009vn}. As a consequence, a large part of MOND theories based on modified Poisson equation and characterized by standard interpolating functions is excluded by Cassini range data (in particular the original TeVeS interpolating function is excluded by our analysis)."
    },
    "1402/1402.5929_arXiv.txt": {
        "abstract": "{We present solutions for the velocity field and mass-loss rates for 2D axisymmetric outflows, as well as for the case of mass accretion through the use of the Lambert W-function. For the case of a rotating radiation-driven wind the velocity field is obtained analytically using a parameterised description of the line acceleration that only depends on radius $r$ at any given latitude $\\theta$. The line acceleration $g(r)$ is obtained from Monte-Carlo multi-line radiative transfer calculations. The critical/sonic point of our equation of motion varies with latitude $\\theta$. Furthermore, an approximate analytical solution for the supersonic flow of a rotating wind is derived, which is found to closely resemble the exact solution. For the simultaneous solution of the mass-loss rate and velocity field, we use the iterative method of our 1D method extended to the non-spherical 2D case. We apply the new theoretical expressions with our iterative method to the stellar wind from a differentially rotating 40\\,$M_{\\sun}$ O5--V main sequence star as well as to a 60\\,$M_{\\sun}$ O--giant star, and we compare our results to previous studies that are extensions of the Castor et al. (1975, ApJ, 195, 157) CAK formalism. Next, we account for the effects of oblateness and gravity darkening. Our numerical results predict an equatorial decrease of the mass-loss rate, which would imply that (surface-averaged) total mass-loss rates are lower than for the spherical 1D case, in contradiction to the Maeder \\& Meynet (2000, A\\&{}A, 361, 159) formalism that is oftentimes employed in stellar evolution calculations for rotating massive stars. To clarify the situation in nature we discuss observational tests to constrain the shapes of large-scale 2D stellar winds. ",
        "introduction": "In recent years considerable progress has been made in our theoretical modelling of rotating massive stars (Maeder \\& Meynet \\cite{maeder:meynet:2012}; Langer \\cite{langer}) and our basic understanding of spherical radiation-driven winds (Puls et al. \\cite{puls:vink}). However, in order to get a grasp on the non-spherical 2D outflows of rotating massive stars, involving B[e] supergiants (Zickgraf et al. \\cite{zickgraf:wolf}), classical Be stars (Porter \\& Rivinius \\cite{porter:rivinius}) as well as luminous blue variable (LBV) outflows (Groh et al. \\cite{groh:hillier}), it is paramount to combine the intrinsically 2D nature of rotation and mass loss (e.g. Lovekin \\cite{lovekin}; Espinosa \\& Rieutord \\cite{espinosa:rieutord}). This is not only required for a basic understanding of massive star evolution, but also for linking the oftentimes non-spherical supernova (SN) data with their progenitors (e.g. Maund et al. \\cite{maund:wheeler}, Hoffman et al. \\cite{hoffman:leonard}).  We need to develop 2D wind models in order to obtain a physical understanding of how rotation might affect both the strength and latitudinal dependence of their outflows. In turn winds may be able to remove significant quantities of angular momentum, potentially down to masses as low as 10-15\\,$M_{\\sun}$ (Vink et al. \\cite{vink:brott}). Whether the mass loss originates from the pole or the equator remains currently unknown. Yet, is of paramount importance for understanding whether rapid rotation is maintained or leads to stellar spin loss (Meynet \\& Maeder \\cite{meynet:maeder}), highly relevant for our understanding of the progenitor evolution of long-duration gamma ray bursts (GRBs).  Previous models of the winds from rotating stars have mostly been 1D models of the equatorial flow versus the polar flow, although one 2D numerical calculation has been performed by Poe (\\cite{poe}). The 1D model of Friend \\& Abbott (\\cite{friend:abbott}, hereafter FA) concerned the influence of stellar rotation on the hydrodynamics of a stellar wind, involving a solution of the fluid equations in the equatorial plane, which included centrifugal forces. They used a form of the radiation force after Castor et al.~(\\cite{castor:abbott:klein}, hereafter CAK), but corrected for the finite angular size of the stellar ``disk'' (see also Pauldrach et al. \\cite{pauldrach:al}).  Bjorkman \\& Cassinelli (\\cite{bjorkm:cassi}, hereafter BC) provided an analytical approximation for the axisymmetric 2D supersonic solution (i.e. for the velocity field and density distribution) of a rotating radiation-driven wind, obtained from the FA 1D model of the equatorial flow. In order to find the streamline trajectories, they rotated the FA 1D solution (in the equatorial plane) up to the initial co-latitude $\\theta_{0}$ of the streamline at the stellar surface, adjusting the equatorial rotation velocity of the central star $v_{\\rm rot}$ by \\mbox{$v_{\\rm rot}\\sin\\theta_{0}$}. The supersonic solutions obtained this way provided the velocity and density as a function of $\\theta_{0}$ and radius $r$, i.e.~in a non-explicit form: given a location $(r,\\theta)$, one needs to find $\\theta_{0}$ of the streamline that passes through that location, by iteratively solving additional equations. As a result, they explained how rotation can lead to the production of a dense equatorial disk around, e.g. Be stars, by means of their wind-compressed disk (WCD) model.  Refinements of the BC model have been also made for simulating the density structure of rotating O-star winds (e.g.~Petrenz \\& Puls \\cite{petrenz:puls}). Owocki et al. (\\cite{owocki:cranmer}) showed that the inclusion of non-radial line forces leads to a small polarwards $v_{\\theta}$ component, which may inhibit disk formation in Be stars. Moreover, Maeder (\\cite{maeder}) showed that gravity darkening as a result of the Von Zeipel (\\cite{zeipel}) effect will generally lead to a polar wind. It should be noted that the hydrodynamical wind models of Owocki and colleagues that lead to polar winds employ an approximated line driving formula from CAK for 1D.  A generally prolate wind structure was however confirmed by the sectorial 1.5D wind models of Pelupessy et al. (\\cite{pelupessy:lamers:vink}) that employed 1D detailed Monte Carlo line acceleration computations of Vink et al. (\\cite{vink:dekoter:lamers}). In their models for B[e] supergiants, Pelupessy et al. also showed that when models are in close proximity to the bi-stability jump it is possible to overcome the polar enhancement due to the Von Zeipel effect, and drive equatorial enhancements, as originally suggested by Lamers \\& Pauldrach (\\cite{lamers:pauldrach}) for Be and B[e] supergiant disk formation. Cure et al. (\\cite{cure:rial}) and Madura et al. (\\cite{madura:owocki}) derived 1D hydrodynamical models for very rapid rotators (above 75\\% of the critical rate) finding a slow solution to the classical CAK theory, which may enable disk formation in Be stars and B[e] supergiants, when accounting for the wind bi-stability effects of Vink et al. (\\cite{vink:dekoter:lamers}). However, again, these models employ a simplified treatment of the line acceleration due to the 2 (or 3) parameter power-law approximation due to CAK. What is eventually required in order to resolve the intricate problem of stellar rotation with mass loss are 3D Monte Carlo radiative transfer calculations in combination with a full hydrodynamic solution. Most published models have necessarily made significant assumptions with respect to either the line-force calculations or the wind hydrodynamics.  We suggested a new parametrisation of the line acceleration (M\\\"uller \\& Vink \\cite{mueller:vink}, hereafter Paper I), expressing it as a function of radius rather than of the velocity gradient as in CAK theory. The implementation of this formalism allowed for local dynamical consistency as we were able to determine the momentum transfer at each location in the wind through the use of Monte Carlo simulations. In Muijres et al. (\\cite{muijres:al}), we tested our hydrodynamic wind solutions and velocity laws by additional explicit numerical integrations of our fluid equation, also accounting for a temperature stratification. These results were in excellent agreement with both our full and our approximated solutions from Paper I. We here build on those results, now deriving analytical solutions for the 2D case concerning both the velocity and density structure in an axisymmetric mass outflow (or inflow) scenario. Furthermore, we extend our iterative method from Paper I for the simultaneous solution of the mass-loss rate and velocity field to the 2D case of a rotating non-spherical stellar wind.  We obtain the velocity field fully analytically without any previous fits to numerical solutions of the fluid equation of FA, if we neglect the polar velocity $v_{\\theta}$ in our model. We are justified in doing so as long as the stellar rotation speeds are well below the critical value where disks are formed, that is, for example for O-star winds, where non-spherical outflows have only been detected in a small ($\\sim20\\%$) minority only involving ``special'' O-type sub-groups (Oe, Onfp) from linear spectropolarimetry observations (Harries et al. \\cite{harries:howarth}; Vink et al. \\cite{vink:davies}).  However, the non-spherical problem for the flow in the equatorial plane including the case of the outflow at the pole (where $\\mbox{$v_{\\theta}\\equiv{}0$}$ for symmetrical reasons) are as well fully analytically solved by our improved 2D wind model without any restrictions to the velocity components and the rotational speed of the central star. For the specific case of a non-spherical radiation-driven wind, we do not rely on the CAK expression for the radiation force, rather we describe the line acceleration as a function of stellar radius $g(r,\\theta{})$ at a given constant co-latitude $\\theta$. In addition, the critical point of our stellar wind is the sonic point (depending on latitude) and not the CAK critical point. The calculation of $g(r,\\theta{})$ is performed through Monte Carlo simulations accounting for multi-line transfer, and the wind parameters are solved simultaneously -- in an iterative way -- for each latitude of interest.  The set-up of the paper is as follows. In Sects.~\\ref{sect-basic-hydro-eq}--\\ref{sect-line-acc-term}, the hydrodynamic equations for a non-spherical axisymmetric steady flow are introduced including a derivation of the mathematical description of the radiative line acceleration as a function of radius for the case of a rotating radiation-driven wind. The process for obtaining the fully analytical 2D solutions is described and discussed in Sect.~\\ref{sect-analyt-sol-Eq-mot}. Here, the velocity field for the entire family of solutions is provided in an explicit general expression from which the solutions for a rotating radiation-driven wind or mass accretion flux (e.g. collapsing protostellar cloud) follow as unique trans-sonic solutions through the critical point. Moreover, an approximate analytical solution for the supersonic flow is presented. In Sect.~\\ref{sect-num-meth}, we describe our numerical computation obtaining the radiative acceleration in our stellar wind models. Furthermore, our iterative method for the determination of the consistent solution for the mass-loss rate in case of a spherical wind is being extended and applied to the wind from a rotating star. In Sect.~\\ref{sect-applic-ostar}, we present the application of our models to a differentially rotating stellar wind from a typical  40\\,$M_{\\sun}$ O5--V-star, including the effects of oblateness and gravity darkening, and from a rotating 60\\,$M_{\\sun}$ O--giant star. We discuss the results, before we summarise and discuss our findings in Sect.~\\ref{sect-disc}. ",
        "conclusions": "\\label{sect-disc}  \\subsection{Comparison with previous models}  The earliest works on CAK-type mass-loss predictions including the effects of stellar rotation via the effective mass in the equation of motion (but without considering gravity darkening) such as those by FA and Pauldrach et al. (\\cite{pauldrach:al}), Petrenz \\& Puls (\\cite{petrenz:puls:2000}) resulted in moderate mass-loss rate enhancements, which were subsequently included in evolutionary calculations (Langer \\cite{langer}). Our results, with mass-loss rates increases by at most 0.1 dex, are in good agreement with these earlier works based on the CAK theory.  When gravity darkening is included, previous studies, such as those of Owocki et al. (\\cite{owocki:cranmer}), Maeder (\\cite{maeder}) and Pelupessy et al. (\\cite{pelupessy:lamers:vink}) all found relatively modest changes for low rotational velocities, but they all favour polar ejection for high rotation rates in close proximity to the break-up limit ($\\Omega$ = 1). The most-oft used formulae for mass-loss enhancement on the basis of the CAK theory is the one given by Maeder \\& Meynet (\\cite{maeder:meynet}). Their famous equation (4.29) reads:  \\begin{equation} \\frac{\\dot{M}(\\Omega)}{\\dot{M}(\\Omega=0)} = \\frac{(1-\\Gamma)^{\\frac{1}{\\alpha}-1}}{\\left[ 1-\\frac{\\Omega^2}{2\\pi G \\rho_{\\rm m}} - \\Gamma \\right]^{\\frac{1}{\\alpha}-1}}, \\label{EqMdotRot} \\end{equation} where $\\Gamma=L/L_{\\rm Edd}=\\kappa L / (4\\pi cGM)$ is the Eddington factor (with $\\kappa$ the total opacity), and $\\alpha$ the $T_{\\rm eff}-$dependent CAK force multiplier parameter. Whilst it is claimed that the mass-loss increase is due to rotation ($\\Omega$) the more relevant reason for the high mass-loss rate is in fact the high $\\Gamma$ factor (see more recent computations by Vink \\cite{vink}, Gr\\\"{a}fener \\& Hamann \\cite{graefener:hamann}, Vink et al. \\cite{vink:muijres}).  Surprisingly, our 2D results that take gravity darkening into account predict a equatorial decrease in the mass-loss rate. This would also imply that the total (cf. Eq.~\\ref{total-mass-loss}) mass-loss rate is now lower than for the spherical 1D case. This is in contradiction to the above Maeder \\& Meynet equation that is oftentimes used in stellar evolution calculations for rotating massive stars.   \\subsection{Comparison with observations}  There are two complementary types of astronomical techniques that can provide meaningful geometric constraints on large-scale 2D wind structures. These are interferometry and linear spectropolarimetry. The current status is that whilst interferometry has been very successful in obtaining results on objects such as classical Be stars, B[e] supergiants, and LBVs, results for O stars are -- to the best of our knowledge --  not yet available due to their smaller radii.  Linear spectropolarimetry can be employed with at least two levels of complexity. The first involves detailed line-profile predictions that can be used to determine the geometric and kinematic environments around young and massive stars (Harries \\cite{harries}, Vink et al. \\cite{vink:harries}), as well as the polarimetric agents (Kuhn et al. \\cite{kuhn:berdyugina}). The lower-level complexity application of linear spectropolarimetry is adopted here, with the emission line reasonably assumed to be unpolarised, but with the continuum responsible for the observed polarisation level: depolarisation (Poeckert \\& Marlborough \\cite{poeckert:marlborough}, Brown \\& McLean \\cite{brown:mclean}). This latter approach is sufficient for our current comparisons. In this case a pole-to-equator density contrast of 5 might lead to a continuum polarisation of $\\simeq$1\\% according to the expectations of Harries et al. (\\cite{harries:hillier}).  If we assume that our most realistic density-contrast predictions concern those models that include gravity darkening, we would anticipate the vast majority of mostly moderately rotating O--dwarfs to be unpolarised, as the density contrast between the pole and equator is close to unity for rotation velocities up to 300 km/s. These overall findings are in good agreement with the linear H$\\alpha$ polarimetry survey of Vink et al. (\\cite{vink:davies}).  An exception might involve the sub-group of Oe stars, for which Vink et al. (\\cite{vink:davies}) list the highest rotation velocities, with $v$sin$i$ values approaching $\\simeq$400 km/s. As this concerns a lower limit (due to the inclination effects), the true rotational velocities of Oe stars might go all the way up to the break-up limit ($\\Omega = 1$). Interestingly, the Oe star HD\\,45314 was indeed shown to have a line effect, but the overall incidence of line effects in Oe stars was low: 1/6. On the basis of the general spectroscopic similarity between Oe and Be stars, and the presence of double-peaked line profiles many astronomers would expect Oe stars to be embedded in circumstellar disks, although it should be noted that spectroscopy alone cannot provide geometric constraints, without complementary imaging, interferometry, or polarimetry, because geometric information from spectral line profiles is not unique, due to the convolution that underlies spectral line formation. In particular, double-peaked profiles can even be noted in spherical shells (e.g. PP).  In the current paper the surprising result is that of polar outflows, whilst it should be noted that linear spectropolarimetry cannot yet distinguish between a prolate (polar) or oblate (equatorial) wind structure (e.g. Wood et al. \\cite{wood:bjorkman}). It will be challenging to reconcile our axisymmetric density-contrast predictions with the general expectation that Oe/Be stars are embedded within equatorial structures. Clearly there are many puzzles remaining in the geometry and formation of circumstellar media around rotating OB stars, and quite possibly additional physical effects, such as magnetic fields, pulsations, and companions may need to be considered.   \\subsection{Summary and Future Work}   In Paper I (M\\\"uller \\& Vink \\cite{mueller:vink}), we proposed a new parametrisation of the radiative line acceleration, expressing it as a function of radius rather than of the velocity gradient (as in CAK theory) which generalised the classical thermal Parker (\\cite{parker}) wind and allowed solving it analytically. This has the advantage of higher accuracy and enables to check an influence of numerical viscosity if used as a test example for numerical simulations. In addition, the implementation of this formalism allows for local dynamical consistency as we are able to determine the momentum transfer at each location in the wind through the use of Monte Carlo simulations.  In this paper, we present a generalization of our model of spherical line-driven winds (in Paper I) to the case of axially symmetric rotating stars. We extend on the results of Paper~I, deriving analytical solutions for the 2D case in an axisymmetric mass outflow or inflow scenario. The generalization approximates the solution of the two-dimensional problem of a rotating wind as a one-parametric set (parameterised by the latitude) of one-dimensional solutions. Here, the separation of the radial motion for individual latitudes is achieved by the basic assumption that the meridional flow velocity is negligible. Assuming furthermore only central external forces, then imply the conservation of angular momentum and leads to an additional centrifugal term in the equation of radial motion.\\footnote{ The centrifugal term was also included by Friend \\& Castor (\\cite{friend:castor}) in an opposite approximation of an extreme transport of angular momentum able to ensure purely radial motion in a co-rotating frame.} We also extend our iterative method for the simultaneous solution of the mass-loss rate and velocity field to the case of a rotating non-spherical stellar wind, using the parameterised description of the line acceleration that only depends on radius -- at any given latitude.  Furthermore, an approximate analytical solution for the supersonic flow of a rotating wind is derived, which closely resembles the exact solution, and we apply the new expressions with our iterative method to (i) the stellar wind from a differentially rotating 40\\,$M_{\\sun}$ O5--V main sequence star, as well as to (ii) a 60\\,$M_{\\sun}$ O--giant star. We subsequently account for the effects of stellar oblateness and gravity darkening for the O5--V main sequence star. The results show an equatorial decrease in the mass-loss rate, which would imply a total mass-loss rate that is lower than for the spherical 1D case, in contradiction to the oft-used Maeder \\& Meynet (\\cite{maeder:meynet}) formalism used in most current stellar evolution calculations for rotating massive stars.  In the future we plan to extend our method to B supergiants, in particular to tackle the problem of disk-formation in B[e] supergiants (Lamers \\& Pauldrach \\cite{lamers:pauldrach}, Pelupessy et al. \\cite{pelupessy:lamers:vink}, Cure et al. \\cite{cure:rial}, Madura et al. \\cite{madura:owocki})."
    },
    "1402/1402.0021_arXiv.txt": {
        "abstract": "{The mid- and far-infrared view on high-mass star formation, in particular with the results from the {{\\it Herschel}} space observatory, has shed light on many aspects of massive star formation. However, these continuum studies lack kinematic information.} {We study the kinematics of the molecular gas in high-mass star-forming regions.} {We complemented the PACS and SPIRE far-infrared data of 16 high-mass star-forming regions from the {{\\it Herschel}} key project EPoS with \\nhp molecular line data from the MOPRA and Nobeyama 45m telescope. Using the full \\nhp hyperfine structure, we produced column density, velocity, and linewidth maps. These were correlated with PACS 70\\micron images and PACS point sources. In addition, we searched for velocity gradients. } {For several regions, the data suggest that the linewidth on the scale of clumps is dominated by outflows or unresolved velocity gradients. IRDC\\,18454 and G11.11 show two velocity components along several lines of sight. We find that all regions with a diameter larger than 1\\,pc show either velocity gradients or fragment into independent structures with distinct velocities. The velocity profiles of three regions with a smooth gradient are consistent with gas flows along the filament, suggesting accretion flows onto the densest regions.} {We show that the kinematics of several regions have a significant and complex velocity structure. For three filaments, we suggest that gas flows toward the more massive clumps are present. % } ",
        "introduction": "Despite their rarity, high-mass stars are important for all fields of astronomy. Within the Milky Way they shape and regulate the formation of clusters, influence the chemistry of the interstellar medium, and may even have affected the formation of the solar system \\citep{Gritschneder2012}. On larger scales, emission from high-mass stars dominates the emission detected from external galaxies. In addition, massive stars are the origin of heavy elements on all scales. Nevertheless, high-mass star formation is far from being understood \\citep{Beuther2007, Zinnecker2007}.  Sensitive infrared (IR) and (sub-)\\,mm Galactic plane surveys together with results from the {\\it Herschel\\footnote{{\\it Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}} space observatory \\citep{Pilbratt2010} have shed new light on the cradles of massive stars/clusters and their early formation. \\citet{Perault1996} and \\citet{Egan1998} discovered extinction patches in the bright mid-IR (MIR) background using the ISO \\citep{Kessler1996} and MSX \\citep{Egan2003} satellites. These extinction patches are similar to the dark patches reported by \\citet{Barnard1919}, which are today known to be connected to low-mass star formation. Soon after this, \\citet{Carey1998} established that the so-called infrared dark clouds (IRDCs) are the precursors of high-mass star formation. Today, the Spitzer observatory Galactic plane surveys GLIMPSE at 3.6\\micron, 4.5\\micron, 5.8\\micron, and 8\\micron \\citep{Benjamin2003} and MIPSGAL at 24\\micron \\citep{Carey2009} allow the systematic search for IRDCs with unprecedented sensitivity \\citep[e.g.][]{Peretto2009}.  While Spitzer improved our MIR view of the Galaxy, the {\\it Herschel} satellite allows observations of the far-IR (FIR). With the PACS \\citep{Poglitsch2010} and SPIRE \\citep{Griffin2010} photometer, high sensitivity and spatial resolution observations between 70\\micron and 500\\micron are possible. From correlating data at MIR through submillimeter and millimeter wavelengths, the picture emerged that most star-forming regions are filamentary \\citep{Andre2010, Menshchikov2010, Molinari2010, Schneider2010, Hill2011, Hennemann2012, Peretto2012}.  In numerical studies, the formation of dense cores and clumps is explained by two scenarios. On the one hand, molecular clouds fragment in a self-similar cascade down to the typical size of dense, quasi-static cores supported by turbulence. These will then form single or multiple gravitationally bound objects \\citep{McKee2003, Zinnecker2007}. On the other hand, in the dynamical theory molecular clouds are formed from large-scale flows of atomic gas as transient objects \\citep{MacLow2004, Klessen2005, Heitsch2008, Clark2012}. Within these transient structures, supersonic turbulence compresses some fraction of the gas to filaments, clumps, and dense cores. If gravity dominates, the cores collapse. In contrast to the quasi-static cores, these cores constantly grow in mass. When, by chance, some cores accrete mass faster than others due to their higher initial gravitational potential, this is called competitive accretion \\citep{Bonnell2004}. Also dynamical, but of reversed reasoning, in the fragmentation-induced starvation scenario by \\citet{Peters2010}, individual massive dense cores build from the large-scale flows, and an accompanying cluster of smaller cores drag away material from the main core and diminish its mass accretion.   The Earliest Phases of Star formation (EPoS, PI O. Krause) is a Guaranteed Time {\\it Herschel} Key Program for investigating 14 low-mass and 45 high-mass star-forming regions. % The low-mass observations have been summarized in \\citet{Launhardt2013}, and the high-mass part has been described in \\citet{Ragan2012}. % The high-mass part of the project provides an excellent target list for studying the kinematics in star-forming regions. This is the ultimate goal of this paper, using \\nhp molecular line data. ",
        "conclusions": ""
    },
    "1402/1402.4091_arXiv.txt": {
        "abstract": "{Precision measurements of the Sunyaev-Zel'dovich effect in clusters of galaxies require excellent rejection of common-mode signals and wide frequency coverage.}{We describe an imaging, efficient, differential Fourier transform spectrometer (FTS), optimized for measurements of faint brightness gradients at millimeter wavelengths.}{Our instrument is based on a Martin-Puplett interferometer (MPI) configuration. We combined two MPIs working synchronously to use the whole input power. In our implementation the observed sky field is divided into two halves along the meridian, and each half-field corresponds to one of the two input ports of the MPI. In this way, each detector in the FTS focal planes measures the difference in brightness between two sky pixels, symmetrically located with respect to the meridian. Exploiting the high common-mode rejection of the MPI, we can measure low sky brightness gradients over a high isotropic background.} {The instrument works in the range $\\sim$ 1$-$20 cm$^{-1}$ (30$-$600 GHz), has a maximum spectral resolution $1/(2 \\ OPD) = 0.063 \\ cm^{-1}$ (1.9 GHz), and an unvignetted throughput of 2.3 cm$^2$sr. It occupies a volume of 0.7$\\times$0.7$\\times$0.33 m$^3$ and has a weight of 70 kg. This design can be implemented as a cryogenic unit to be used in space, as well as a room-temperature unit working at the focus of suborbital and ground-based mm-wave telescopes. The first in-flight test of the instrument is with the OLIMPO experiment on a stratospheric balloon; a larger implementation is being prepared for the Sardinia radio telescope.}{}} ",
        "introduction": "About 1\\% of the photons of the cosmic microwave background (CMB) that cross a rich cluster of galaxies undergo inverse Compton scattering against the hot electrons of the intracluster plasma. This phenomenon is called the thermal Sunyaev-Zel'dovich effect (SZE) and is a powerful tool for astrophysical and cosmological investigation (Sunyaev and Zeldovich \\cite{SunZel70}, (Sunyaev and Zeldovich \\cite{SunZel72}, Rephaeli \\cite{reph95}, Birkinshaw \\cite{birk99}, Carlstrom et al. \\cite{Carl02}).  Most of the power of the thermal SZE is due to its very characteristic spectrum: in the direction of a cluster of galaxies, the SZE induces a {\\sl decrease} of the brightness of the CMB at frequencies below 217 GHz, and an {\\sl increase} of the brightness of the CMB at frequencies above 217 GHz. In the nonrelativistic limit, the spectrum of the SZE can be computed from the Kompaneets equation (Kompaneets \\cite{Komp57}) and only depends on the Comptonization Parameter $y$, \\begin{equation} y=\\int n_{e}\\sigma_{T}\\frac{k_{B}T_{e}}{m_{e}c^{2}} d\\ell \\ , \\end{equation} where the integral is along the line of sight. Using the $y$ parameter, the brightness variation $\\Delta I_{t}$ with respect to the average cosmic microwave brightness $I$ can be expressed as \\begin{equation} \\frac{\\Delta I_{t}}{I}= y\\frac{x^{4}e^{x}}{(e^{x}-1)^{2}}\\bigl[ x  coth(x/2)-4 \\bigr] \\ .\\label{eq:tsze} \\end{equation} Typical values of $y$ are $10^{-4} - 10^{-5}$, so the SZE signal $\\Delta I_t$ is much weaker than the average brightness $I$ of the CMB. However, the SZE spectrum does not depend on the distance (redshift) of the cluster of galaxies, so clusters become important cosmological probes. Moreover, the SZE spectrum depends linearly on both the temperature and density of the intracluster plasma. Information from the SZE significantly complements information obtained from X-ray brightness (proportional to the density squared) in investigating the physics of the clusters. The thermal SZE is only the most evident interaction between the CMB and the plasma of the cluster. The peculiar motion of the cluster induces the kinematic SZE, and the nonthermal component of the plasma produces additional distortions (Wright \\cite{Wrig79}).  Our ability of measuring the SZE has recently reached maturity with the surveys of large, dedicated ground-based telescopes (see e.g. Reichardt et al. \\cite{Reic13}, Hasselfield et al. \\cite{Hass13} and references therein) at low frequencies, and with the space-based survey of the Planck satellite (see Planck collaboration \\cite{Plan13} and references therein), which covers the whole spectral range where the SZE is present.  All these are multiband photometric measurements. The number of independent bands useful to this purpose is limited to a few for ground-based observations, and was 6$-$7 in the case of the Planck satellite. The parameters that describe the cluster and the foregrounds can be more than seven: in principle, more bands and/or complementary observations in other spectral regions are needed.  Spectroscopic measurements can help very much in this situation. They allow us to extract the weak SZE signal even in the presence of important foreground emission, and to estimate the parameters of the cluster by removing the degeneracy of photometric measurements (see de Bernardis et al. \\cite{debe12} and section \\ref{ps} for a specific discussion).  \\begin{figure}[ht] \\centering \\includegraphics[width=8.5cm]{fig2.pdf}\\\\ \\caption{{\\bf Bottom:} Thermal (continuous line) and kinematic (dashed line) SZE for $y=5 \\times 10^{-4}$, $v_{p}=1000 km/s$ and $T_{e}=8.2keV$, in MJy/sr.  {\\bf Top:} residual atmospheric emission at stratospheric balloon altitude (continuous line, in MJy/sr as well). Four sub-bands (see text) suitable for balloon-borne obserations of the SZE are also plotted in the same panel (dashed lines). }\\label{SZE} \\end{figure}  Here we present an instrument optimized for spectroscopic measurements of the SZE. Owing to the continuous nature of the spectrum (see Fig. \\ref{SZE}), high spectral resolution is not a design-driver. A coarse spectral resolution ($\\sim$ 0.5 cm$^{-1}$) is more than enough to provide a sufficient number of independent spectral data. Better resolution could be useful to assess whether there is contamination from molecular lines, but this is important only for low galactic latitudes. Instead, the instrument is required to be imaging (feeding a detector array in its focal plane), to have a wide frequency coverage (almost two decades at mm/submm wavelengths), and to be differential, that is, to measure the difference in brightness between the cluster direction and the surrounding directions, rejecting the bright isotropic background from the CMB. The class of the Fourier transform spectrometers (FTSs), and in particular the Martin-Puplett interferometer (MPI, Martin and Puplett \\cite{Mart70}), fulfills all these requirements.  Fourier transform spectrometers are imaging instruments with high spectral resolution and efficiency in the IR, with a variety of applications in chemistry, remote sensing, atmospheric studies, planetary exploration and IR astronomy. Notable space-borne implementations of differential FTSs for astrophysics and cosmology are the Far InfraRed Absolute Spectrometer (FIRAS) on-board of the COBE satellite (Mather et al. \\cite{Math90}, Fixsen et al. \\cite{Fixs94}) and the Spectral and Photometric Imaging Receiver (SPIRE)  FTS on-board of the Herschel satellite (Griffin et al. \\cite{Grif10}). These two compare  the brightness of the sky with the brightness produced by an internal blackbody reference.  In our instrument the two input ports instead collect radiation from two independent sky fields, so that each detector in the focal plane of the spectrometer measures the brightness difference between two. In this way the common-mode signals from the instrument, the residual atmosphere and the isotropic CMB component, are all rejected, and only brightness gradients are measured. For this reason we use the achronym DFTS (differential Fourier transform spectrometer) for our instrument. A previous implementation of this configuration at shorter wavelengths is the Bear interferometer at the Canada France Hawaii Telescope (CFHT) (see e.g. Maillard \\cite{Mail95}).  First light for our DFTS will be on a stratospheric balloon flight of the OLIMPO experiment (Masi et al. \\cite{Masi03}); additional configurations have been studied for the Spectroscopic Active Galaxies And Clusters Explorer (SAGACE) satellite proposal (de Bernardis et al. \\cite{debe10}), for the millimetron satellite mission (Smirnov et al. \\cite{Smir12}), and for the Sardinia radio telescope (SRT, Grueff et al. \\cite{Grue04}).  The paper is organized as follows: in section \\ref{ps} we show that adding spectral measurements of the SZE to photometric measurements results in a significant enhancement of the ability to separate degenerate parameters that describe the SZE and foreground contamination.  In section  \\ref{efficient} we quantitatively describe the operation of the DFTS and demostrate its efficiency. In section  \\ref{design} we describe the optical design strategy for the instrument. In section \\ref{mech} we describe the optical optimization  and the opto-mechanical design. Then (section \\ref{prob}) we discuss the most critical problems related to this configuration.  In section \\ref{test} we describe the test and commissioning of the instrument.  We conclude in section \\ref{disc}. ",
        "conclusions": "\\label{disc}  We have developed a differential Fourier-transform spectrometer optimized for low-medium-resolution spectroscopy of faint brightness gradients. The instrument is efficient, features large throughput ($\\sim  2.3 \\ cm^2 sr$) and medium resolution ($\\sim 1.9$ GHz FWHM, independent of frequency), in a relatively compact setup. The instrument is suitable for suborbital operation and will be tested in the forthcoming flight of the OLIMPO stratospheric balloon.  We tested several possible systematic effects. The only one resulting in a potentially non-negligible systematics is the $CMMR$, for which we set an upper limit $CMMR < 0.2\\%$. For measurements of the SZ effect with OLIMPO, where the DFTS works at room temperature, the common-mode background on the detectors is due to the instrument itself (mainly the room-temperature mirrors and the WGs), the residual atmosphere, the CMB, and interstellar dust. The resulting common-mode brightness, multiplied by our CMRR upper limit, is about (6, 20, 60, 160) MJy/sr in our four bands, to be compared to SZ signals about (-0.2, 0.2, 0.7, 0.8) MJy/sr. An offset removal observation strategy is evidently in order. The SZ measurement will be performed by subtracting from a first observation of the brightness difference between the cluster and the reference direction, a second observation of the brightness difference between two other nearby blank-sky directions. With a stability of the common-mode emissions better than 0.05\\%, this strategy allows good measurements of the SZ. We stress that what we just described is the worst case allowed by our upper limit on the CMRR. Our instrument model suggests a much better situation. The confirmation will come from in-flight measurements.  A second implementation of the DFTS was designed for operation in the W-band at the Sardinia radio telescope.  With a different mechanism for the mirror motion (see e.g. Schillaci et al. \\cite{Schi11} and references therein), the system can be implemented for operation in a cryogenic environment, suitable for space-based operation. This is the case of the SAGACE (de Bernardis et al. \\cite{debe10}) and millimetron (Smirnov et al. \\cite{Smir12}) missions. Here the CMRR offset due to instrument emission will be reduced by at least a factor 100, resulting in very competitive spectrometers for wide-band continuum measurements."
    },
    "1402/1402.1482_arXiv.txt": {
        "abstract": "{We report the discovery of the planets WASP-20b and WASP-28b along with measurements of their sky-projected orbital obliquities. WASP-20b is an inflated, Saturn-mass planet (0.31\\,\\mjup; 1.46\\,\\rjup) in a 4.9-day, near-aligned ($\\lambda = 8.1 \\pm 3.6 \\degr$) orbit around CD-24~102 ($V$=10.7; F9). WASP-28b is an inflated, Jupiter-mass planet (0.91\\,\\mjup; 1.21\\,\\rjup) in a 3.4-day, near-aligned ($\\lambda = 8 \\pm 18 \\degr$) orbit around a $V$=12, F8 star. As intermediate-mass planets in short orbits around aged, cool stars ($7^{+2}_{-1}$ Gyr for WASP-20 and $5^{+3}_{-2}$ Gyr for WASP-28; both with \\teff\\ $<$ 6250\\,K), their orbital alignment is consistent with the hypothesis that close-in giant planets are scattered into eccentric orbits with random alignments, which are then circularised and aligned with their stars' spins via tidal dissipation. } ",
        "introduction": "Planets that transit relatively bright host stars ($V$$<$13) are proving a rich source of information for the nascent field of exoplanetology. To date, the main discoverers of these systems are the ground-based transit surveys HATNet and SuperWASP and the space mission Kepler \\citep{2004PASP..116..266B, 2006PASP..118.1407P, 2010Sci...327..977B}.  One parameter that, uniquely, we can determine for transiting planets is obliquity ($\\Psi$), the angle between a star's rotation axis and a planet's orbital axis. We do this by taking spectra of a star during transit: as the planet obscures a portion of the rotating star it causes a distortion of the observed stellar line profile, which manifests as an anomalous radial-velocity (RV) signature known as the Rossiter-McLaughlin (RM) effect \\citep{1924ApJ....60...15R,1924ApJ....60...22M}. The shape of the RM effect is sensitive to the path a planet takes across the disc of a star relative to the stellar spin axis. If we have a constraint on the inclination of the stellar spin axis relative to the sky plane ($I_*$), then we can determine the true obliquity ($\\Psi$), otherwise we can only determine the sky-projected obliquity ($\\lambda$). The two are related by:  $\\cos \\Psi = \\cos I_* \\cos i_{\\rm P} + \\sin I_* \\sin i_{\\rm P} \\cos \\lambda$  \\noindent where $i_{\\rm P}$ is the inclination of the orbital axis to the sky plane.  The obliquity of a short-period, giant planet may be indicative of the manner in which it arrived in its current orbit from farther out, where it presumably formed. As the angular momenta of a star and its planet-forming disc both derive from that of their parent molecular cloud, stellar spin and planetary orbital axes are expected to be, at least initially, aligned. Migration via tidal interaction with the gas disc is expected to preserve initial spin-orbit alignment \\citep{1996Natur.380..606L,2009ApJ...705.1575M}, but almost half (33 of 74) of the orbits so far measured are misaligned and approximately 10 of those are retrograde.\\footnote{Ren\\'e Heller maintains a list of measurements and references at \\url{http://www.physics.mcmaster.ca/~rheller/content/main_HRM.html}}  These results are consistent with the hypothesis that some or all close-in giant planets arrive in their orbits by planet-planet and/or star-planet scattering, which can drive planets into eccentric, misaligned orbits, and tidal friction, which circularises, shortens and aligns orbits (see the following empirical-based papers: \\citealt{2010A&A...524A..25T,2010ApJ...718L.145W,2011Natur.473..187N,2012ApJ...757...18A}; also see the following model-based papers: \\citealt{2007ApJ...669.1298F,2008ApJ...678..498N,2010ApJ...725.1995M, 2011Natur.473..187N}).  Systems with short tidal timescales (those with short scaled orbital major semi-axes, $a$/\\rstar, and high planet-to-star mass ratios) tend to be aligned (\\citealt{2012ApJ...757...18A} and references therein). A broad range of obliquities is observed for stars with \\teff$>$6250\\,K, for which tidal realignment processes may be inefficient due to the absence of a substantial convective envelope \\citep{2010ApJ...718L.145W,2010ApJ...719..602S}. Limiting focus to stars more massive than 1.2\\,\\msol, for which age determinations are more reliable, \\citet{2011A&A...534L...6T} noted that systems older than 2.5~Gyr tend to be aligned; this could be indicative of the time required for orbital alignment or of the timescale over which hotter stars develop a substantial convective envelope as they evolve.  A major hurdle to overcome for any hypothesis involving realignment is tidal dissipation seems to cause both orbital decay and realignment on similar timescales \\citep{2009MNRAS.395.2268B}. Alternative hypotheses suggest that misalignments arise via reorientations of either the disc or the stellar spin and that migration then proceeds via planet-disc interactions (e.g. \\citealt{2010MNRAS.401.1505B,2011MNRAS.412.2790L,2012ApJ...758L...6R}). However, observations of debris discs well aligned with their stellar equators suggest that tilting of the star or the disc rarely occurs \\citep{2014MNRAS.438L..31G}.  Here we present the discovery and obliquity determinations of the two new transiting planets WASP-20b (CD-24~102) and WASP-28b (2MASS~J23342787$-$0134482), and interpret the results under the hypothesis of migration via scattering and tidal dissipation. ",
        "conclusions": "We present the discovery of WASP-20b, a Saturn-mass planet in a 4.9-day orbit around CD-24~102, and WASP-28b, a Jupiter-mass planet in a 3.4-day orbit around a $V$=12 star. Based on their masses, orbital distances, irradiation levels and metallicities, the radii of the planets (1.46\\,\\rjup\\ for WASP-20b and 1.21\\,\\rjup\\ for WASP-28b) are consistent with the predictions of the empirical relations of \\citet{2012A&A...540A..99E}: 1.33\\,\\rjup\\ for WASP-20b and 1.26\\,\\rjup\\ for WASP-28b.  We find both planets to be in prograde orbits and, as the \\vsini\\ values are consistent with expectations for $v$ for such stars, it is probable that $I_*\\approx 90\\degr$ and so $\\Psi \\approx \\lambda$. Assuming the spin axis of WASP-20 to be in the sky plane, the orbit of WASP-20b is aligned or near-aligned with the stellar spin ($\\lambda = 8.1 \\pm 3.6 ^\\circ$). Although our determination for WASP-28b ($\\lambda = 8 \\pm 18 ^\\circ$) is consistent with an aligned orbit, its low impact parameter results in a greater uncertainty.  Both host stars are near the posited boundary between efficient and inefficient aligners of \\teff\\ $\\approx$ 6250\\,K \\citep{2010ApJ...718L.145W,2010ApJ...719..602S}. Both systems, with their low obliquity determinations, are consistent with the observation that systems with cool host stars and short expected tidal timescales are aligned (\\citealt{2012ApJ...757...18A} and references therein). With \\mplanet/\\mstar\\ = 0.00025 and $a$/\\rstar\\ = 9.3, WASP-20b experiences relatively weak tidal forces. Its relative tidal dissipation timescale of $\\tau_{\\rm CE} = 2.6 \\times 10^{13}$ yr is the longest for the cool-host group barring HD\\,17156\\,b (see section 5.3 of \\citealt{2012ApJ...757...18A}). \\footnote{Note that these timescales are relative and that \\citet{2012ApJ...757...18A} plot the timescales divided by $5 \\times 10^9$.} With \\mplanet/\\mstar\\ = 0.00084 and $a$/\\rstar\\ = 9.5, $\\tau_{\\rm CE} = 2.5 \\times 10^{12}$ yr for WASP-28b, placing it within the main grouping of aligned, cool-host systems. Assuming the planets did not arrive in their current orbits recently, the advanced age of both systems, $7^{+2}_{-1}$ Gyr for WASP-20 and $5^{+3}_{-2}$ Gyr for WASP-28, mean tidal dissipation would have occurred over a long timescale.  The planets may have arrived in their current orbits via planet-disc migration \\citep{1996Natur.380..606L,2009ApJ...705.1575M} or via the scattering and circularisation route \\citep{2007ApJ...669.1298F,2008ApJ...678..498N,2010ApJ...725.1995M, 2011Natur.473..187N}. Assuming the latter, the planets may have scattered into aligned orbits or they may have aligned with the spins of their host stars via tidal interaction, though how this could have occurred without the planets falling into their stars is currently a mystery \\citep{2009MNRAS.395.2268B}."
    },
    "1402/1402.6698_arXiv.txt": {
        "abstract": "The successful miniaturization of extremely accurate atomic clocks invites prospects for satellite missions to perform precise timing experiments. This will allow effects predicted by general relativity to be detected in Earth's gravitational field. In this paper we introduce a convenient formalism for studying these effects, and compute the fractional timing differences generated by them for the orbit of a satellite capable of accurate time transfer to a terrestrial receiving station on Earth, as proposed by planned missions. We find that (1) Schwarzschild perturbations will be measurable through their effects both on the orbit and on the signal propagation, (2) frame-dragging of the orbit will be readily measurable, and (3) in optimistic scenarios, the spin-squared metric effects may be measurable for the first time ever.  Our estimates suggest that a clock with a fractional timing inaccuracy of $10^{-16}$ on a highly eccentric Earth orbit will measure all these effects, while for a low Earth circular orbit like that of the Atomic Clock Ensemble in Space mission, detection will be more challenging. ",
        "introduction": "Atomic clock technology has improved dramatically over the past decade. Present-day Earthbound atomic clocks boast fractional timing inaccuracies of $\\sim10^{-18}$ \\cite{Chou, Middelmann, Hinkley}. The most stable space-qualified clock, with stability $\\sim10^{-16}$ has been built for the Atomic Clock Ensemble in Space (ACES) mission, and will be placed on the International Space Station in 2016. As we show in this paper, a variety of relativistic effects are extremely important at these accuracy levels. Optical time and frequency transfer over free space has reached a residual stability of $\\sim10^{-18}$ over distances of a few kilometers \\cite{Giorgetta}. As such technology improves, it is expected that highly accurate clocks on Earth will be commonly used for orbit determination.  The theory of general relativity explains gravitation as a geometrical phenomenon arising from a curved four-dimensional spacetime. The Earth, carrying mass and momentum, determines the trajectories - both spatially and temporally - of satellites which fall freely around it. The Einstein equivalence principle tells us that regardless of the position or velocity of a freely falling satellite, its onboard clock will tick routinely in equal intervals of time in its frame. This paper simulates an experiment to test gravitational physics by capturing this notion: a satellite carries an ultraprecise atomic clock and broadcasts tick signals to an Earth-based receiving station, whose arrival times are compared with a local clock.  \\begin{figure}[!h] \\begin{center} \\includegraphics[scale =0.6,trim=2cm 10cm 10cm 1.2cm,clip=true]{setup.pdf} \\caption{\\label{fig:orientation} Schematic of the system for an eccentric orbit. Here we define the angles in Earth's spin reference frame. The two black points mark the satellite orbit's perigee and the Earth receiving station. Two angles appear negative, as they are defined positively when measured from apogee. For computational simplicity, and without loss of generality, the Earth spin axis is always along the $z$ direction.} \\end{center} \\end{figure}  A direct consequence of the geometrical description of gravity is that a clock in a shallow gravitational field runs more rapidly than one wallowing more deeply in the field. For Earth satellites, this affects the redshift at the $G M_\\oplus/(c^2 r_\\oplus) \\sim 10^{-9}$ level (see Fig. \\ref{fig:all_signals}). Thus a fractional timing stability of $10^{-16}$ implies a sensitivity to gravitational time dilation of $10^{-7}$. We show that at these remarkable accuracy levels, higher-order relativistic effects must be taken into account.  An experiment with a freely falling satellite, which periodically sends tick signals to an Earth receiving station, is sensitive to the full four-dimensional trajectory. This is in contrast to Earth-based fixed clock experiments, which probe gravity purely through its effect on time, and not on space. State-of-the-art space-based atomic clocks, and modern time transfer technology promise orbit tracking to unprecedented accuracy. As the accuracy of clocks increases, a host of relativistic effects become important.  In this paper we address the forward problem: calculating the relativistic observables from Earth-orbiting clocks. Our calculations suggest that the Shapiro delay, Mercury-like orbit precession, frame-dragging, and possibly even spin-squared effects, will be detectable by future satellite timing missions. Our approach is general and applicable to a multitude of orbits around any gravitating body. An interactive computer program which calculates all the effects discussed in this paper is available as an online supplement, downloadable from the journal website.  As examples, we will consider an eccentric orbit, as well as a low Earth circular orbit. An eccentric orbit is particularly advantageous because it brings the clock through various gravitational field strengths and thus permits the relativistic effects to modulate. Also, we find that the higher perigee velocities of the elliptical orbit boost the relativistic orbit effects by about an order of magnitude, but do not affect relativistic signal transmission effects. ",
        "conclusions": ""
    },
    "1402/1402.3597_arXiv.txt": {
        "abstract": "{ Although accounting only for a small fraction of the baryonic mass, the dust has a profound impact on the physical processes at play in galaxies. Thus, to understand the evolution of galaxies, it is essential not only to characterize dust properties per se, but also in relation to global galaxy properties. To do so, we derive the dust properties of the galaxies of a volume limited, K-band selected sample, the \\textit{Herschel} Reference Survey (HRS).  We gather infrared photometric data from 8\\microns\\ to 500\\microns\\ from \\textit{Spitzer}, WISE, IRAS and \\textit{Herschel} for all of the HRS galaxies. Draine \\& Li (2007) models are fit to the data from which the stellar contribution has been carefully removed. We find that our photometric coverage is sufficient to constrain all of the parameters of Draine \\& Li (2007) models and that a strong constraint on the 20-60\\microns\\ range is mandatory to estimate the relative contribution of the photo-dissociation regions to the infrared spectral energy distribution (SED). The SED models tend to systematically under-estimate the observed 500\\microns\\ flux densities, especially for low mass systems. We provide the output parameters for all of the galaxies, i.e. the minimum intensity of the interstellar radiation field, the fraction of polycyclic aromatic hydrocarbon (PAH), the relative contribution of PDR and evolved stellar population to the dust heating, the dust mass and the infrared luminosity. For a subsample of gas-rich galaxies, we analyze the relations between these parameters and the main integrated properties of galaxies, such as stellar mass, star formation rate, infrared luminosity, metallicity, H$\\alpha$ and H-band surface brightness, and the far-ultraviolet attenuation. A good correlation between the fraction of PAH and the metallicity is found implying a weakening of the PAH emission in galaxies with low metallicities and, thus, low stellar masses. The intensity of the diffuse interstellar radiation field and the H-band and H$\\alpha$ surface brightnesses are correlated, suggesting that the diffuse dust component is heated by both the young stars in star forming regions and the diffuse evolved population.  We use these results to provide a new set of infrared templates calibrated with \\textit{Herschel} observations on nearby galaxies and a mean SED template to provide the $z$=0 reference for cosmological studies. For the same purpose, we put our sample on the $SFR$-$M_*$ diagram. The templates are compared to the most popular infrared SED libraries, enlightening a large discrepancy between all of them in the 20-100\\microns\\ range. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.5008_arXiv.txt": {
        "abstract": "We carry out non-rotating high-resolution calculations of the solar global convection, which resolve convective scales of less than 10 Mm. To cope with the low Mach number conditions in the lower convection zone, we use the reduced speed of sound technique (RSST), which is simple to implement and requires only local communication in the parallel computation. In addition, the RSST allows us to expand the computational domain upward to about $0.99 R_{\\odot}$ as it can also handle compressible flows. Using this approach, we study the solar convection zone on the global scale, including small-scale near-surface convection. In particular, we investigate the influence of the top boundary condition on the convective structure throughout the convection zone as well as on small-scale dynamo action. Our main conclusions are: 1. The small-scale downflows generated in the near-surface layer penetrate into deeper layers to some extent and excite small-scale turbulence in the region $>0.9R_\\odot$, where $R_\\odot$ is the solar radius. 2. In the deeper convection zone ($<0.9R_\\odot$), the convection is not influenced by the location of the upper boundary. 3. Using an LES approach, we can achieve small-scale dynamo action and maintain a field of about $0.15-0.25 B_\\mathrm{eq}$ throughout the convection zone, where $B_\\mathrm{eq}$ is the equipartition magnetic field to the kinetic energy. 4. The overall dynamo efficiency varies significantly in the convection zone as a consequence of the downward directed Poynting flux and the depth variation of the intrinsic convective scales. ",
        "introduction": "The solar convection zone is filled with turbulent thermal convection due to its superadiabatic stratification, which is maintained by the radiative energy flux from the core of the Sun. Due to the influence of rotation, the anisotropic convection transports angular momentum as well as energy. This maintains mean flows such as the meridional flow and the differential rotation. These flows are thought to be important for the dynamo and related eleven-year cyclic variations of the sunspot numbers \\citep{1955ApJ...122..293P,1995A&A...303L..29C,1999ApJ...518..508D, 2010ApJ...709.1009H,2010ApJ...714L.308H}. In addition, the turbulent thermal convection itself is also important for the generation of the magnetic field. Idealized simulations \\citep[e.g.,][]{1996JFM...306..325B,1999ApJ...515L..39C} and the realistic calculations for the photosphere \\citep{2007A&A...465L..43V,0004-637X-714-2-1606} have suggested that chaotic motions have the ability to maintain magnetic energy through a small-scale dynamo process. Such dynamos have not been studied so far on the the global scale, i.e., the entire convection zone. \\par The purpose of this study is to achieve multi-scale convection, including scales from $10$~Mm (smaller than supergranulation) up to $\\sim 200$~Mm (global scale). We accomplish this by approaching the solar surface up to $0.99R_{\\odot}$ using unprecedented resolution, studying in particular the influence of the location of the upper boundary on the convective structure in the deeper parts of the convection zone. In addition, we use our approach to study the transport and the generation of  the magnetic field by turbulent convection in the absence of rotation, i.e., we study the operation of a small-scale dynamo in the highly stratified convection zone. \\par The paper is organized as follows: We provide an overview of the reduced speed of sound technique in \\S \\ref{S:rsst}, introduce our numerical setting in \\S \\ref{model}, and show the results in \\S \\ref{result},present the properties based on the obtained spatial distributions in \\S \\ref{structure}, discuss the energy balance using the integrated flux in \\S \\ref{integrate_flux}, analyze , the properties of the convection in cases without the magnetic fields using the spherical harmonic expansion in \\S \\ref{sh}, and analyze the cases with magnetic fields using spherical harmonic expansion and the probability density function in \\S \\ref{sh_mag}, investigate the transportation and the generation of the magnetic field in the convection zone in \\S \\ref{local_dynamo}, and summarize our investigation in \\S \\ref{discussion}. ",
        "conclusions": "\\label{discussion} We carry out the high-resolution calculations of the solar global convection which resolve  10 Mm-scale convection smaller than the supergranulation using the RSST (described in \\S \\ref{S:rsst}). The RSST leads to a simple algorithm and requires only local communication in the parallel computing. In addition, this method has the capability to access the real solar surface without undermining the important physics. This enables us to capture near-surface small-scale convection while keeping a global domain. Our main conclusions are given as follows: 1. Small-scale convection is excited close to the surface ($>0.9R_\\odot$), when we expand our domain upward to $0.99\\,R_{\\odot}$ to capture the near-surface layers with small pressure scale heights. 2. In deeper convection zones ($<0.9R_\\odot$), the convection flow is not influenced by the location of the top boundary and the assumed thickness of the thermal boundary layer. Changing the position of the top boundary or thickness of the cooling layer does not lead to significant differences in the convective structure and properties of the local dynamo, in terms of the power-spectrum, the probability density function and the local dynamo. 3. Using an LES approach we can achieve small-scale dynamo action and maintain a field of about $0.15-0.25 B_\\mathrm{eq}$ throughout the convection zone. 4. The overall dynamo efficiency varies significantly in the convection zone as a consequence of the downward directed Poynting flux and the depth variation of the intrinsic convective scales. For a fixed numerical resolution, the dynamo-relevant scales are better resolved in the deeper convection zone and are therefore less affected by numerical diffusivity, i.e., the effective Reynolds numbers are larger. \\par Recently, \\cite{2012PNAS..10911928H} suggested that there is a significant difference between the calculation by \\cite{2008ApJ...673..557M} and flow velocities inferred from local helioseismology. The observed convection amplitude is much smaller than that in the hydrodynamic calculation. We expected that the modification to the location of the upper boundary may help to resolve this issue. Conclusion 1 above, however, suggests that the difference between the calculation and the observation becomes larger when we move the top boundary further upward. We see, nevertheless, some indication that the kinetic energy is reduced on the largest scales when a local dynamo is present. At the same time, we cannot fully rule out the possibility that unresolved small-scale turbulence near the top boundary contributes to this issue, since the intrinsic scale of the convection near our top boundary is still not well resolved. \\par Conclusion 2 is one of the most important results in this study, since it suggests that previous calculations such as by \\cite{2008ApJ...673..557M} are physically reasonable in the deeper convection zone even if the top boundary condition is placed significantly below the solar surface. \\par We summarize our simulation results in a schematic picture shown in Fig. \\ref{local_dynamo_fig}. The conclusion that the dynamo is ineffective near the top comes from the comparison of the kinematic and saturated spectra and the only moderate reduction of the effective shear in the upper convection zone (see \\S \\ref{local_dynamo}). Since several aspects, in particular with regard to the local dynamo, require higher resolution before being directly applicable to the solar convection zone, Fig. \\ref{local_dynamo_fig} does not necessarily apply to the real Sun. High resolution simulations of a local dynamo in the solar photosphere \\citep{2007A&A...465L..43V} suggest that efficient dynamo action is possible even in the presence of the pumping effect. These simulations use, however, a grid resolution of about a factor of $100$ larger than our setup, which is currently inapplicable for global scale convection simulations. Therefore, the dynamo RMS field strength of $0.15-0.25 B_\\mathrm{eq}$ can likely be considered a lower limit. \\par One issue we cannot address in this study is the problem regarding the small magnetic Prandtl number ($\\mathrm{Pm}\\sim 10^{-3}$) in the solar convection zone, since we adopt numerical diffusivities which assume that the magnetic Prandtl number is around unity. Several authors argued that smaller magnetic Prandtl numbers make local dynamos less efficient \\citep[e.g.][]{2004PhRvL..92e4502S,2004PhRvL..92n4501B}. In other words, a small magnetic Prandtl number requires a larger magnetic Reynolds number in order to achieve a super-critical dynamo. While this can be a significant problem for numerical simulations with rather moderate Reynolds numbers, this is less likely an issue in the solar convection zone with Reynolds numbers as large as $\\mathrm{Rm}\\sim 10^{11}$ \\citep{2011ApJ...741...92B}. Thus we believe that our approach relying only on numerical diffusivities can capture the physics of the local dynamo in the solar convection zone. \\par Rotation is not included in this study although it plays a key role in organizing large-scale flows and enabling large-scale dynamo action. Studying the near-surface small-scale convection in a setup with rotation is of particular interest to us for investigation of the origin of the near-surface shear layer observed on the Sun. Our investigation of this is a work-in-progress and will be presented in a forthcoming paper. Further progress can be made by using non-uniform grids, such as nested grids or adaptive mesh refinement, which are useful in order to overcome the significant differences in spatial and temporal dynamic scales. These methods are already implemented in our numerical code \\citep{2012A&A...548A..74H}."
    },
    "1402/1402.0077_arXiv.txt": {
        "abstract": "{Stellar flybys in star clusters are suspected to affect the orbital architecture of planetary systems causing eccentricity excitation and orbital misalignment between the planet orbit and the equatorial plane of the star.} {We explore whether the impulsive changes in the orbital elements of planets, caused by an hyperbolic stellar flyby, can be fully damped by the circumstellar disk surrounding the star. The time required to disperse stellar clusters is in fact comparable to circumstellar disk's lifetime. Since we perform 3D simulations, we can also test the inclination excitation and damping.} {We have modelled in 3D with the SPH code VINE a system made of a solar type star surrounded by a low density disk with a giant planet embedded in it approached on a hyperbolic encounter trajectory by a second star, of similar mass and with its own disk. Different inclinations between the disks, planet orbit and star trajectory have been considered to explore various encounter geometries. We focus on an extreme configuration where a very deep stellar flyby perturbs a Jovian planet on an external orbit. This allows to test in full the ability of the disk to erase the effects of the stellar encounter.} {We find that the amount of mass lost by the disk during the stellar flyby is less than in 2D models where a single disk was considered. This is mostly related to the mass exchange between the two disks at the encounter. The damping in eccentricity is slightly faster than in 2D models and it occurs on timescales of the order of a few kyr. During the flyby both the disks are warped due to the mutual interaction and to the stellar gravitational perturbations, but they quickly relax to a new orbital plane. The planet is quickly dragged back within the disk by the tidal interaction with the gas. The only trace of the flyby left in the planet system, after about $10^4$ yr, is only a small misalignment, lower than $9^o$, between the star equatorial plane and the planet orbit. } {In a realistic model based on 3D simulations of star--planet--disk interactions, we find that stellar flybys cannot excite significant eccentricities and inclinations of planets in stellar clusters. The circumstellar disks hosting the planets damp on a short timescale all the step changes in the two orbital parameters produced during any stellar encounter. All records of past encounters are erased.} ",
        "introduction": "\\label{intro}  Stars generally do not form alone but in bound clusters on a timescale of about $1$ Myr \\citep{hille97, palla00}. The embedded cluster phase lasts only a few Myr and, according to \\cite{allen07}, \\cite{fall09}, \\cite{chan10}, and \\cite{duk12}, after about $10$ Myr, $90\\%$ of stars born in clusters have dispersed into the field. This occurs because, after a few crossing times \\citep{allen07}, embedded clusters lose their gaseous component, the system becomes super--virial and the cluster rapidly dissolves. A cluster, before its dissipation,  is a crowded place and its dynamical evolution is characterized by frequent stellar flybys. The close encounters can have significant effects on both disks and planetary systems surrounding the stars. They can alter the disk morphology causing also mass loss and excite planet orbits. It is expected that, in the period of higher frequency of stellar encounters, circumstellar disks are still present around stars. The lifetime of disks is in fact comparable to the cluster bound phase being in the range $3$--$6$ Myr, as suggested by \\cite{hll06}, or even longer, up to $10$--$12$ Myr, as recently argued by \\cite{bell13}. Therefore, to explore the effects of close stellar encounters on planetary systems in clusters, one has to account for the presence of circumstellar disks around the stars.  Recent N--body numerical modeling of the dynamical evolution of planetary systems around stars in clusters \\citep{malm1,malm2,malm3,zaktre} have shown that scattering interactions with other stars  may excite the eccentricity and inclination of planets populating the outer parts of the system. The steep changes in the orbital elements would be observable also at present and might contribute to explain the large eccentricities of extrasolar planets discovered so far and the presence of misaligned systems. However, \\cite{marpi13} have shown that the planet's eccentricity, excited during a close stellar flyby, is damped on a short timescale ($\\sim 10$ kyr) by the interaction with the circumstellar disk harbouring the planet in spite of the disk's low density related to its viscous evolution and photoevaporation occurred during the planet growth. As a consequence, the coincidence of the cluster lifetime with that of the circumstellar disks plays against a significant role of stellar flybys in affecting the present planet dynamical architecture.  The simulations presented in \\cite{marpi13} were based on the 2D code FARGO and considered a single disk around the star with planets while the perturber star was assumed diskless. In this paper we extend that model using a 3D SPH code and including two disks, one around each star. This allows us to better model close stellar flybys where the two disks closely interact. In addition, we can also explore the excitation of the planet inclination and its subsequent damping due to the interaction with the disk \\citep{marne09,klebi11,cres07}. We concentrate in this paper on the effects of a single deep close stellar encounter on a planet still embedded in its birth disk. If the effects of an extremely close flyby can be erased by the disk, we expect that also the cumulative effects of multiple but more distant encounters can be absorbed by the disk. We will also explore the effects on the disk morphology of the stellar tidal force and of the mutual disk interaction during the encounter.  The paper is organized as follows. In Sect. 2 we describe the numerical algorithm used to model the disks and planet evolution during and after a hyperbolic close encounter with another star. In Sect. 3 we outline the different flyby configurations we have modeled while Sect 4,5 and 6 are devoted to the presentation of the numerical outcomes. In Sect. 7 we discuss our results and their implications for the evolution of planetary systems in stellar clusters. ",
        "conclusions": "\\label{results}  The 3D simulations of stellar encounters in the presence of circumstellar disks, presented in this paper, confirm that the orbital changes a planet undergoes during a close stellar flyby are quickly damped by the interaction with the disk. The only observable difference after the close approach may be a residual misalignment between the planet--disk orbital plane and the star equatorial plane. In fact, during the encounter, the disk harbouring the planet is warped by the interaction with the passing star and its disk. Gradually, it relaxes on a new plane slightly tilted respect to the original one. The amount of tilting depends on the geometrical configuration of the system prior the encounter and in our models we find a maximum value of about $9^o$. Subsequently, the planet is pulled back into the disk by the tidal interaction with the gas and its orbit also becomes inclined respect to the star equatorial plane. In case of repeated stellar flybys the final tilting of the planet orbit may be larger than  $9^o$  but it appears unlikely that this mechanism may be responsible for the large planet inclinations found in some extrasolar systems \\citep{triaud10}. The inclination evolution due to a series of stellar flybys would appear as a random walk with jumps that may lead to positive or negative changes of the planet (and disk) inclination. The stepsize of this random walk would also be small, since we consider an extreme case with a very deep flyby. An event like this would be rare in the lifetime of the star within the cluster and its probability would rapidly decrease with time.  We find some significant differences between 2D and 3D simulations. \\begin{itemize} \\item The mass lost by the primary disk during the stellar flyby depends on the 3D geometry of the encounter and does not lead to the large mass loss observed in 2D simulations. The difference is mainly due to the inclusion of a disk around the passing star that exchanges mass with the primary disk during the close encounter and the use of an SPH code that allows us to better track the evolution of the gas far away from the stars. \\item The planet migrates inward at a slower rate in 3D simulations compared to 2D models. The migration of an eccentric orbit, like that of the planet just after the encounter, appears to be significantly less efficient in 3D models. \\item The initial eccentricity excited by the stellar flyby is damped more quickly in 3D models and this is partly due to the reduced amount of mass lost by the disk. \\end{itemize} We also observed a slower eccentricity damping efficiency when the orbit of the planet is inclined respect to the disk due to a slightly weaker disk--planet interaction.  Concerning the effects of stellar flybys on planetary systems in stellar cluster, we do not expect any signature of the stellar encounters in the eccentricity of the planets. Some misalignment between the planet orbit and the star equatorial plane may appear but it would be limited. A period of fast type II migration may occur after the stellar encounter since any gap created by the planet before the flyby is destroyed during the strong interaction between the stars and the disks at the pericenter. A new gap is created when the planet is finally pulled back into the disk midplane and its orbit circularized. The rate at which the disks are destroyed in clusters may be slower than expected since when two stars have an encounter, their disks exchange mass rather than loosing it.  \\cite{marpi13} explored the consequences of stellar encounters on systems of 3 planets and a  similar exploration would be interesting also in 3D to test the occurrence of a ``jumping Jupiter'' chaotic phase and derive its final outcome. The parameter space is really huge and many different configurations may be envisaged in light of recent discoveries of extrasolar planetary systems. However, according to our results, it appears reasonable to expect that if the damping works in the extreme conditions we have simulated (low disk density, very deep encounter) it should be effective also in different, potentially more complex, scenarios."
    },
    "1402/1402.0846_arXiv.txt": {
        "abstract": "We use high dynamic range, high-resolution $L$-band spectroscopy to measure the radial velocity variations of the hot Jupiter in the $\\tau$~Bo\\\"{o}tis planetary system. The detection of an exoplanet by the shift in the stellar spectrum alone provides a measure of the planet's minimum mass, with the true mass degenerate with the unknown orbital inclination. Treating the $\\tau$~Boo system as a high flux ratio double-lined spectroscopic binary permits the direct measurement of the planet's true mass as well as its atmospheric properties.  After removing telluric absorption and cross-correlating with a model planetary spectrum dominated by water opacity, we measure a 6-$\\sigma$ detection of the planet at K$_{p}$ = 111 $\\pm$ 5 km$/$s, with a 1-$\\sigma$ upper limit on the spectroscopic flux ratio of 10$^{-4}$.  This radial velocity leads to a planetary orbital inclination of $i$ = 45~$^{+3}_{-4}$~degrees and a mass of M$_{P}$ = 5.90~$^{+0.35}_{-0.20}$~M$_{\\rm Jup}$.  We report the first detection of water vapor in the atmosphere of a non-transiting hot Jupiter, $\\tau$~Boo\\,b. ",
        "introduction": "Since the first detection of an exoplanet around a main-sequence star \\citep{Mayor95}, astronomers have discovered hundreds of exoplanets using the radial velocity (RV) technique \\citep{Wright12}.  This powerful tool for discovering exoplanets only measures the minimum, or ``indicative\" mass, ${M_{\\rm p} \\sin{i}}$, leaving a degeneracy between two interesting properties of the system, one physical and one orbital.  While the indicative mass is useful for statistical studies, uniquely measuring the true mass of an exoplanet not only yields a key physical property, but also furthers our understanding of planetary formation and evolution via measurement of the true mass distribution of exoplanets  \\citep{Zucker01, Weiss13}. For example, \\citet{HoTurner11} demonstrate that knowledge of the true mass distribution is necessary to convert a minimum mass, ${M_{\\rm p} \\sin{i}}$, into an estimate of a planet's true mass for RV-detected systems.  One class of objects for which the planet mass can be directly determined are those that transit their host star.  Hundreds of transiting planets have been discovered and characterized, and the ongoing \\textit{Kepler} mission has found potential exoplanet candidates numbering in the thousands \\citep{Borucki11, Batalha13}.  Transit events also provide atmospheric information through transmission spectroscopy and secondary eclipses.  Investigators have measured spectra of some of the larger transiting planets \\citep{Charbonneau02}, leading to the discovery of species such as water, methane and carbon monoxide \\citep{Knutson12, Berta12, Crossfield13, Baskin13}.  A variety of spectral retrieval methods have been used to verify the results, confirming the hypothesis that O- and C-bearing gases are present in the atmospheres of hot Jupiters \\citep{Madhusudhan09, Line12}.  Recently, a technique previously used to detect low mass ratio, spectroscopic binary stars has been applied to stars known to host exoplanet systems.  The direct RV detection of an exoplanet involves separating the planetary and stellar components spectroscopically.  The high flux ratio between the primary and companion makes these detections difficult, but not impossible, thanks to modern infrared echelle spectrographs.  \\citet{Snellen10} detected CO on HD\\,209458\\,b with a precision of 2 km/s using the R=100,000 Very Large Telescope CRIRES instrument.  This detection provided the spectroscopic orbit of the system and determined the true mass of the planet.  %  Indeed, this technique can be applied to non-transiting, RV-detected exoplanets to extract the unknown inclination and true mass.  CRIRES was also used to detect CO on $\\tau$~Bo\\\"{o}tis b \\citep{Brogi12}, the first ground-based detection of a short-period non-transiting exoplanet atmosphere, a result confirmed shortly thereafter by \\citet{Rodler12}.   These studies provide the true mass of the planet and probe the chemical composition of its atmosphere.  A combination of high signal-to-noise, high spectral resolution, and coverage of multiple CO overtone lines was required to achieve the sensitivity required for these detections.  Despite the agreement between the two groups, the direct detection of exoplanets, especially $\\tau$~Boo\\,b, has a storied history.  \\citet{CollCam99} first reported a detection of reflected light from $\\tau$~Boo\\,b more than a decade ago and \\citet{Wiedemann01} reported a possible detection of CH$_4$ in the planet's atmosphere soon thereafter, using the same orbital solution.  However, the planetary velocity from the earlier results disagrees with more recent findings discussed above.  Here we report a detection of water vapor in the atmosphere of $\\tau$~Boo\\,b, using spectroscopic observations centered around 3.3~$\\mu$m. The $\\tau$~Bo\\\"{o}tis system is a stellar binary comprised of a F-type `A' component and a M-dwarf `B' component.  The hot Jupiter $\\tau$~Boo\\,b orbits the larger `A' component with a period of 3.312 days and a minimum (indicative) mass M$\\sin{i}$ = 3.87M$_{Jup}$  \\citep{Butler97}. Our detection confirms the mass determined by \\citet{Brogi12} and serves to further characterize the atmospheric chemistry of this exoplanet.  Data from five epochs reveal orbital motion of $\\tau$~Boo\\,b that is consistent with that reported by \\citet{Brogi12}. The planetary template spectrum used in the cross-correlation is dominated by water vapor opacity, providing strong evidence of water in the atmosphere of a non-transiting hot Jupiter for the first time. ",
        "conclusions": "The detection of $\\tau$~Boo\\,b has been a difficult quest.  Contradictory results have been published over the past 15 years and the difficulty in both thoroughly removing the telluric absorption and also extracting the diminutive planetary signal is not to be underestimated.  Multiple observations at different wavebands improve the validity of the detection, and facilitate a variety of statistical tests to ensure an accurate measurement.  Here, the validity of the most likely solutions for the RV of $\\tau$~Boo\\,b are explored.  Of the two prominent velocities that fit the data, one is shown to result from the systematics of the cross-correlation analysis performed on this high flux ratio spectroscopic binary system.  Interestingly, previous velocity studies of the planet in reflected light \\citep{CollCam99} and methane \\citep{Wiedemann01} retrieved values close to this artifact.  We suggest that residual water vapor in the atmosphere, after the telluric removal, might have been responsible for the false value previously reported from infrared observations. As methods for telluric corrections have improved over the past decade, this problem can be overcome and the correct radial velocity retrieved.  Our analysis gives a 6-$\\sigma$ detection of the planet at K$_{p}$ = 111 $\\pm$ 5 km$/$s for $\\tau$~Boo\\,b, with a 1-$\\sigma$ upper-limit on the 3.3~$\\mu$m planet/star spectroscopic flux ratio of 10$^{-4}$.  To determine the significance of our detection, we injected synthetic signals at a variety of planetary velocities and planet-to-star spectroscopic flux ratios and constructed the chi-square surface of the maximum likelihood fits.  The orbital velocity is in good agreement with previous RV amplitude detections via CO by \\citet{Brogi12} and \\citet{Rodler12}, and our analysis reveals the presence of water vapor in the planet's atmosphere.  Furthermore, using a stellar mass of 1.341$^{+0.054}_{-0.039}$ M$_\\odot$ \\citep{Takeda07} and a stellar velocity semi-amplitude of 0.4664 $\\pm$ 0.0033 km$/$s \\citep{Brogi12}, we derive a planetary mass of M$_{P}$ = 5.90~$^{+0.35}_{-0.20}$~M$_{\\rm Jup}$ with an orbital inclination of $i$ = 45~$^{+3}_{-4}$~degrees.  The technique presented here is in its nascent stages, and the work is by no means complete.  Additional quantitative characterization of the physical properties (temperature, opacity) and composition (and thus estimates of vertical mixing) of $\\tau$~Boo\\,b's atmosphere will require both significant simulations of the atmospheric radiative transfer and data analysis along with additional data at longer and shorter wavelengths.  Such work is beyond the scope of this letter, but expanded studies of $\\tau$~Boo using the same methods described above, but also including molecules such as CH$_{4}$ and CO at longer wavelengths than examined here, are underway.  Indeed, although the mole fractions of methane should be insignificant for $\\tau$~Boo\\,b, we have searched for the molecule using the methods described herein and have no detection to report in our data.  With careful analysis in the future, the relative abundances of these molecules in the atmospheres of non-transiting exoplanets can be ascertained.   Further applications of this technique, using water and methane as template molecules, to a variety of additional hot Jupiter exoplanets will be reported elsewhere."
    },
    "1402/1402.5464_arXiv.txt": {
        "abstract": "We have obtained WIYN and SOAR $i^{\\prime}$ images of BL Lacertae objects (BL Lacs) and used these to detect or constrain the flux of the host galaxy. Under common standard candle assumptions these data provide estimates of, or lower bounds on, the redshift. Our targets are a set of flat-spectrum radio counterparts of high flux {\\it Fermi} Large Area Telescope (LAT) sources, with sensitive spectral observations showing them to be continuum-dominated BL Lacs. In this sample 5 of 11 BL Lacs yielded significant host detections, with standard candle redshifts $z=0.13-0.58$. Our estimates and lower bounds are generally in agreement with other redshifts estimates, although our $z=0.374$ estimate for J0543$-$5532 implies a significantly sub-luminous host. ",
        "introduction": "BL Lac objects are an important, extreme sub-class of AGN. The large numbers of BL Lacs being detected in the gamma-rays by the {\\it Fermi} LAT \\citep{2lac} show that these blazars are a dominant contributor to the GeV sky. This flux-limited gamma-ray sample also provides a unique opportunity to probe the evolution of the BL Lac phenomenon over cosmic time \\citep{aet14}.  BL Lac evolution is controversial, with various authors finding negative \\citep[e.g.][]{rec00}, positive \\citep[e.g.][]{mc13} or negligible evolution \\citep[e.g.][]{cac02}. Much of this uncertainty stems from the difficulty in determining redshifts; for many BL Lacs the high jet-associated continuum flux dominates emission from any broad-line region and results in very low equivalent width for galactic host absorption features. Despite intensive spectroscopic campaigns \\citep{shaw13b} this often precludes direct BL Lac redshift measurement, challenging their utility for population and evolution studies \\citep{aet14}. Thus alternate methods of estimating or constraining the source redshift can be quite valuable \\citep{shaw13a}.  In this paper, we image {\\it Fermi} BL Lacs to search for host flux in the wings of the nuclear point spread function (PSF). Our targets are drawn from the `2LAC' \\citep{2lac} LAT blazars for which sensitive spectroscopy, reported in \\citet{shaw13b}, shows them to be continuum-dominated BL Lacs at very high significance, but was not able to determine a redshift. Since hosts are increasingly difficult to detect at high redshift, we excluded BL Lacs whose spectra showed intervening absorbers requiring $z>0.7$ or excluded continuum flux from hosts at lower $z$; this left 226 targets. We observed objects from this list when conditions and gaps in other observing programs allowed.  It has been argued that BL Lac hosts are nearly uniform giant ellipticals with $M_R$ = $-$22.5 \\citep{sba05}. Thus several authors have used such imaging to constrain the source redshift \\citep[e.g.][]{sba05,meisner1,nilet12}. One can also search for the elliptical flux in high S/N spectra; when not detected this provides a lower limit on the redshift \\citep[e.g.][]{sba06,shaw13b,sand13}. Conversely, when direct spectroscopic redshifts measurements are in hand (or later become available), the flux measurements can be used to test the standard candle hypothesis and the evolution of the BL Lac hosts. Here, we follow the method described by \\citet{meisner1}, although we use an updated host magnitude calibration derived from absorption lines features measured in a deep spectroscopic survey \\citep{shaw13a}. ",
        "conclusions": "A primary goal of this exercise is to use the host fluxes to constrain the source redshift, adopting the standard candle hypothesis. To this end we use the Hubble diagram curve computed in \\citet{meisner1}. This curve follows an elliptical host formed at $z=2$ and evolving according to the \\citet{fr97} models to the observed $z$ in our standard cosmology, computing the observed flux by folding through the $i^\\prime$ filter. However, a recent spectral survey of {\\it Fermi} BL Lacs \\citep{shaw13a} finds that the host luminosity of these $\\gamma$-ray sources is 0.4 magnitudes fainter than the $M_R$ = -22.9 $\\pm$ 0.5 reported by \\citet{sba05}. We thus amend the Hubble diagram by normalizing the evolving models to this decreased luminosity at $z=0$.  \\begin{figure} [h!!] \\begin{center} \\includegraphics[width=9.5cm]{fig5.eps} \\end{center} \\caption{$i^\\prime$ Hubble diagram for gamma-ray BL Lac hosts, after \\citet{meisner1}, but normalized to a standard candle magnitude $M_R =-22.5$.} \\end{figure}   Using this curve we translate the host magnitude values and upper and lower statistical and systematic errors to redshift estimates and error ranges. These are reported in Table 2. When only an upper limit on the flux is available, this translated to a redshift lower limit. The redshift estimates for detected hosts varied from 0.13 to 0.58, with lower limits $0.42-0.93$. We consider these limits conservative, in the sense that we use the revised (less luminous) standard candle calibration above. However, insofar as some BL Lacs undoubtedly have substantially sub-luminous hosts, individual sources may indeed appear at lower $z$.  \\begin{table*}[ht!!] \\caption{Fitting Results: Calibrated Magnitudes, Sizes and Redshifts} \\centering \\begin{tabular}{lrrlrrll} Name &SED$^a$&i$^{\\prime}$$_{nucl}$&i$^{\\prime}$$_{Host}$&$f_{host}/f_{nucl}$ &R$_e$ (kpc)&$z_{im}$&$z_{sp}$\\\\ \\hline J0114+1325& H &17.21&20.15$_{-0.037-0.29}^{+0.038+0.41}$  &0.07&8.32$\\pm$0.14&0.583$_{-0.006-0.052}^{+0.007+0.064}$ &$0.61-1.63$\\\\ J0115+2519& H &17.68&18.80$_{-0.015-0.05}^{+0.015+0.05}$  &0.36&12.53$\\pm$0.55&0.358$_{-0.001-0.006}^{+0.001+0.007}$&$0.37-1.63(>0.268)$\\\\ J0222+4302& I &15.13&$>$19.21                        &$<$0.02&-&$>$0.42                                            &$0.12-1.67$\\\\ J0316+0904& H &16.01&18.91$_{-0.033-0.412}^{+0.035+0.672}$&0.07&-&0.372$_{-0.004-0.076}^{+0.005+0.103}$              &$0.12-1.66$\\\\ J0543$-$5532& H &17.10&18.92$_{-0.012-0.17}^{+0.012+0.20}$&0.19&-&0.374$_{-0.002-0.023}^{+0.003+0.027}$             &$0.27-2.57$[=0.273]\\\\ J0558$-$7459& - &17.23&$>$20.52                        &$<$0.05&-&$>$0.66                                            &$0.29-2.20(>0.475)$\\\\ J0700$-$6610& I &16.28&$>$20.08                        &$<$0.03&-&$>$0.57                                            &$0.39-1.92$\\\\ J0721+7120& I &14.01&16.17$_{-0.0024-0.193}^{+0.0038+0.237}$&0.14&9.54$\\pm$0.32&0.127$_{-0.001-0.073}^{+0.002+0.092}$             &$0.14-2.61$\\\\ J1023$-$4336& H &15.26&19.90$_{-0.090-0.689}^{+0.090+2.36}$&0.03&-&$^b$0.534$_{-0.014-0.115}^{+0.019+0.482}$             &$0.34-2.24$\\\\ J1026$-$8543& L &16.88&$>$20.31                        &$<$0.04&-&$>$0.62                                            &$0.32-2.30$\\\\ J1110$-$1835& L &18.69&$>$21.70                        &$<$0.06&-&$>$0.93                                            &$0.51-2.23$\\\\ \\hline \\end{tabular} \\leftline{ Tabulated quantities: Name, SED class, Core magnitude, host magnitude and errors or limit, host radius, } \\leftline{ inferred redshift or limit, spectroscopic constraints on redshift } \\leftline{\\footnotesize $^a$ SED class from \\citet{2lac}, based on the synchrotron peak frequency.} \\leftline{\\footnotesize $^b$ Good statistical, but marginal $\\sim 1\\sigma$ systematic significance. Corresponding lower limit is $z>0.42$} \\end{table*}  \\citet{shaw13a} analyzed high S/N spectra of a large number of {\\it Fermi} BL Lacs. In a method parallel to the present imaging host search, they detected or placed limits on the BL Lac host by measuring the flux of a spectral component from a standard candle elliptical. After appropriate k-correction and slit-loss corrections, these measurements provided host redshift estimates or lower bounds. In addition, sometimes intervening metal line absorption systems were detected. These provide firm, model-independent lower bounds on the host redshift. These spectroscopically allowed ranges and lower bounds are listed in the last column of Table 2; we give here the ranges derived for a standard candle magnitude $M_R$ = -22.5, for consistency with our imaging results.  Since in this program we targeted BL Lacs without known redshift, none of our sources has a spectroscopic $z$ in \\citet{shaw13a}. For the five sources with host detection, two (J0316+0904 and J0543$-$5532) have imaging redshift estimates consistent with the spectra-derived bounds in that paper; the low significance detection of J1023$-$4336 is also consistent. Three (J0114+1325, J0115+2519, and J0721+7120) lie at slightly lower $z$ than the spectral lower bound, but are within 3$\\sigma$, and J0115+2519 is fully consistent with the strict lower limit provided by an intervening absorber. However, since completing this study \\citet{pita13} have published new high quality VLT/X-shooter spectra of a number of blazars, including J0543$-$5532, measured here. For this source, they infer $z=0.237$ based on a weak Ca II H/K doublet and a Na I absorption line. This redshift, near the lower bound of  \\citet{shaw13a} is well below our imaging estimate, implying a host that is substantially fainter than our standard candle assumption. At this $z$ our measured host flux implies an absolute magnitude $M_R = -21.67^{+0.01+0.20}_{-0.01-0.17}$, $0.8$mag ($1.7\\sigma$) away from our assumed standard candle luminosity. The host absolute magnitude is consistent with, but more accurate than, the spectroscopically estimated value in \\citet{pita13}.  In six cases we derive lower bounds on the redshift; J1023$-$4336 may be interpreted as a lower bound of $z> 0.42$.  These limits are always more constraining than those extracted from the spectroscopic study. For J0558$-$7459 our new bound is also stronger than that obtained from the intervening absorption line system. Thus these bounds may be useful in BL Lac population studies \\citep[e.g.][]{aet14}.  With median $f_{host}/f_{nucl} = 0.07$ our sources are very strongly dominated by the non-thermal nuclear core flux. This is in contrast to the HST study of \\citet{scarpa00}, where of 69 BL Lacs with resolved hosts, 37 had $f_{host}/f_{nucl}>1$. Our large core dominance is similar to that found in \\citet{meisner1} and, as noted there, it may be attributed to the fact that these are gamma-ray selected BL Lacs and thus should have high alignment  between the jet axis and the Earth line-of-sight, increasing the core dominance. In addition, these sources were drawn from the BL Lacs lacking redshifts even after extensive spectroscopy with 8\\,m-class facilities \\citep{shaw13a,shaw13b}. Since sources with brighter hosts allow easier absorption line redshift measurements, the remainder (including the sources studied here) should have an especially high core dominance.  In previous studies \\citep{scarpa00,urry00,meisner1} the BL Lacs were seen to have an excess of faint galactic companions, indicating that they were located in cluster environments and that they may have recent interaction activity. In general, the relatively poor seeing and modest image depth achieved during this project prevented us from identifying very faint companions. For 2 objects, we did find bright companion galaxies nearby. J0114+1325 had two companions within a 3$^{\\prime\\prime}$ ($\\sim 20$\\,kpc) radius. These companions were close enough that they overlap significantly with the BL Lac host wings, making accurate flux measurement difficult; we found their i$^{\\prime}$ magnitudes to be roughly 20.5 and 20.8. J0222+4302 had 4 surrounding bright companions located 11$-12^{\\prime\\prime}$ ($\\sim 70$\\,kpc) from the core, with i$^{\\prime}$ magnitudes of 18.9, 19.3, 19.5, and 20.9.  \\bigskip \\bigskip  We have detected hosts in the images of 5/11 BL Lacs, plus one marginal detection. Assuming a standard candle host luminosity these provided redshift estimates $z=0.127 - 0.583$ ($z_{med}$ = 0.37). The minimum redshifts for the remainder are also larger than available spectroscopic lower limits and so are useful, as well. Perhaps unsurprisingly, our host detections were for four HBLs (high spectral peak energy, relatively low luminosity) sources and one intermediate peak source. The other intermediate peak sources and the two LBLs (low peak, higher power) sources in our sample yielded only lower limits on $z$; these include the highest limit, $z>0.93$ in this sample.  When no other redshift estimate is available, these imaging-derived values can be useful for statistical purposes, e.g. in population studies. However, as emphasized by the spectroscopic $z$ recently derived for J0543$-$5532, the standard candle hypothesis is only statistically useful, at best, and should be subject to further study. Indeed \\citet{pita13} estimate that two of their BL Lac hosts have $M_R < -24$, even further from the expected standard candle value. Our host flux measurements thus remain useful whenever an independent redshift is derived. The prospects for further spectroscopic redshifts of our imaged hosts are, in fact, good; these are excellent targets for spatially resolved spectroscopy, especially with Integral Field Unit (IFU) feeds, which can isolate host spectral features from the wings of the BL Lac."
    },
    "1402/1402.2591_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the key frontiers in modern cosmology is measuring the temperature and polarization anisotropies of the Cosmic Microwave Background (CMB) radiation. The anisotropies encode a wealth of information about inflation~\\cite{baumann09}, re-ionization~\\cite{zaldarriaga08}, and the large-scale structure of the universe \\cite{smith08}. From the discovery of the CMB in 1965 by Penzias and Wilson~\\cite{penzias65} through to the Planck satellite of today \\cite{planck13a}, measurements of the CMB have been a key ingredient in establishing the standard model of cosmology.  The CMB is relic light from the early universe. The temperature of the universe before $t\\sim380,000$ years after the Big Bang was high enough that the universe was filled with a plasma of coupled photons and charged particles. As the universe expanded and cooled, neutral hydrogen formed, capturing the free electrons. This event, called recombination, allowed the photons to travel without scattering electromagnetically. This free-streaming background of photons is the CMB that we observe today. Structure in the universe from that era is imprinted on the CMB as hot and cold spots, which means the CMB is the oldest direct picture of the universe we have. The CMB is sometimes called the ``smoking gun'' of the Hot Big Bang scenario.  The uniformity of the CMB temperature across the sky actually presents a challenge to the Hot Big Bang scenario as it was originally conceived. The largest volume of the universe that could have been in causal contact at the time of recombination has an apparent angular size of today of $\\sim2^\\circ$. However, the CMB temperature is highly uniform across the entire sky, to better than one part in $10^4$. This paradox is called the Horizon Problem. Inflation, conceived by Guth \\cite{guth81} and independently by Starobinsky \\cite{starobinsky80}, postulates that there was an early time in the universe's history when the universe expanded exponentially, growing a small causally-connected volume in the pre-inflation era into a large enough volume to explain the uniformity of the CMB. Quantum fluctuations in the field driving the inflationary expansion seeded density perturbations that later grew through gravitational collapse into the fluctuations seen in the CMB, and much later grew into the large-scale structure of the universe today. The quantum fluctuations also perturbed the rest of the metric to create a background of gravitational waves that left an imprint on the polarization of the CMB. Precision measurements of the CMB temperature and polarization are therefore a powerful probe to learn about the physics of inflation. ",
        "conclusions": ""
    },
    "1402/1402.2072_arXiv.txt": {
        "abstract": "We consider the possibility of formation of cosmic strings with black holes as beads. We focus on the simplest setup where two black holes are formed on a long cosmic string. It turns out the in absence of a background magnetic field and for observationally viable values for cosmic string tensions, $\\mu<2\\times 10^{-7}$, the tension of the strut in between the black holes  has to be less than the  ones that run into infinity. This result does not change if a cosmological constant is present. However if the background magnetic field is turned on, we can have stable setups where the tensions of all cosmic strings are equal. We derive the equilibrium conditions in each of these setups depending on whether the black holes are extremal or non-extremal. We obtain cosmologically acceptable solutions with  solar mass black holes and intragalactic strength cosmic magnatic field. ",
        "introduction": "In absence of a background electromagnetic field, two black holes uniformly accelerating from each other are described by the Lorentzian section of the $C-$metric \\cite{Kinnersley:1970zw}. Since there is no background field to provide the acceleration, the metric in general has either conical deficits from the black hole to infinity (cosmic strings), or a conical surplus (cosmic rod)  in between the black holes. In the first case, the Euclidean section of the $C-$metric then describes a tunneling process in which a cosmic string is snapped and two black holes are created at the near ends \\cite{Hawking:1995zn}.  In presence of a background electromagnetic field, black hole pairs accelerating away from each other are described by the Ernst metric \\cite{Ernst1976}. The background magnetic field provides the force that accelerates the (oppositely charged) black holes away from each other. The Euclidean section again describes the process of black hole pair creation in a background magnetic field  \\cite{Garfinkle:1993xk}. This is due to the fact that the negative potential energy of the created pair of black holes in the electromagnetic field is balanced by their rest mass energy.  In \\cite{Emparan:1995je}, the probability of black hole pair creation with a string running between them was found to be suppressed in comparison with the case in which there is no strings. Pairs of black holes can be also created in presence of cosmological constant and during inflation too \\cite{Mann:1995vb,Plebanski:1976gy,Bousso:1996au}.  In all the above papers, it was assumed that only one end of any cosmic string present is attached to  a black holes. We would like to see if configurations in which strings are attached to both ends of the black holes are possible. We are not interested in calculating the instanton solutions  representing the probability of black hole pair creation with strings attached to the  ends; as the probability is often suppressed by the area of the black holes,  cosmic strings attached to them  will reduce the probability further. Instead, we are interested in calculating the classical equilibrium of the configurations irrespective of the mechanism that leading to their formation.  We begin by considering the basic formalism that describes two black holes threaded by a straight cosmic string of infinite length. We find a stable configuration only when the tension of the cosmic string segments running from the black holes to infinity is larger than the segment between the black holes. We then examine generalizations of this configuration in the presence of a cosmological constant and of a constant magnetic field, and obtain necessary conditions for stable configurations.  We find  extremal and near-extremal  stable black hole solutions for cosmologically viable values of the intergalactic magnetic field. ",
        "conclusions": "Using C-metric and its extensions, with the cosmological constant or the magnetic field, we obtained solutions that describe two black holes are on a long cosmic string. We showed that, without the magnetic field, the tension of the strut between the black holes has to be less than the tension of the outer strings running to infinity. Only in the presence of a  background magnetic field is it  possible to have black holes on a cosmic string with uniform tension.  Although the parameters must be appropriately adjusted, cosmologically viable solutions with solar mass black holes exist for intragalactic strength cosmic magnetic field strength. It would be interesting so see if more detailed numerical solutions for such black hole beads could be constructed in which the string is represented by a vortex \\cite{oai:arXiv.org:gr-qc/9506054} so as to better study the phenomenological prospects of these kinds of configurations.   The equilibrium of the above configuration in presence of a cosmic magnetic field is independent of the tension of the cosmic string and only relies on the black hole mass and charge parameters. A large massive black hole can be stabilized by a small tension cosmic string in the presence of an appropriate magnetic field. Thus the lensing \\cite{oai:arXiv.org:astro-ph/9702033} and gravitational waves emitted from cusps on such strings \\cite{oai:arXiv.org:1002.0652} could be in principle be enhanced by the presence of black holes even if the tension of strings is below the detectability limit of probes. One can also investigates the effect of such configuration on the two point correlation of the inflaton \\cite{ARXIV:0908.0543}. This gives further motivation to investigate these signatures of the setup in future."
    },
    "1402/1402.4242_arXiv.txt": {
        "abstract": "Within the nonlinear relativistic mean field model, we show that both the pressure of symmetric nuclear matter at supra-saturation densities and the maximum mass of neutron stars are sensitive to the skewness coefficient $J_0$ of symmetric nuclear matter. Using experimental constraints on the pressure of symmetric nuclear matter at supra-saturation densities from flow data in heavy ion collisions and the astrophysical observation of a large mass neutron star PSR J0348+0432, with the former favoring a smaller $J_0$ while the latter a larger $J_0$, we extract a constraint of $-494 \\mathrm{MeV}\\leq J_0\\leq -10 \\mathrm{MeV}$. This constraint is compared with the results obtained in other analyses. ",
        "introduction": "Determination of the equation of state (EOS) of asymmetric nuclear matter (ANM) is one of fundamental questions in contemporary nuclear physics and astrophysics. The exact knowledge on the EOS of ANM provides important information on the in-medium nuclear effective interactions which play a central role in understanding the structure and decay properties of finite nuclei as well as the related dynamical problems in nuclear reactions~\\cite{Bla80,LiBA98,Dan02,Bar05,Ste05,Che07a,LCK08,Tsa12,Che14}. The EOS of ANM also plays a decisive role in understanding a number of important issues in astrophysics including the structure and evolution of neutron stars as well as the mechanism of supernova explosion~\\cite{Gle00,Lat04,Lat12,SNRev,Hem12,Mei13}. Conventionally, the EOS of ANM is given by the binding energy per nucleon as functions of nucleon density $\\rho$ and isospin asymmetry $\\delta$, i.e., $E(\\rho, \\delta)$, and some bulk characteristic parameters defined at the saturation density $\\rho_0$ of symmetric nuclear matter (SNM) are usually introduced to quantitatively characterize the EOS of ANM. For example, the energy $E_0(\\rho_0)$ and incompressibility $K_0$ of SNM as well as the symmetry energy $E_{\\mathrm{sym}}(\\rho_0)$ and its slope parameter $L $ are the four famous lower-order bulk characteristic parameters of EOS of ANM. These bulk parameters defined at $\\rho_0 $ provide important information on both sub- and supra-saturation density behaviors of the EOS of ANM~\\cite{Che09,Che11a}.  Based on the empirical liquid-drop-like model analyses of high precision data about nuclear masses, the $E_0(\\rho_0)$ is well known to be about $-16$ MeV. The incompressibility has been determined to be $K_0 = 230 \\pm 20$ MeV from analyzing experimental data of nuclear giant monopole resonances (GMR)~\\cite{Bla80,You99,Shl06,Col09,Pie10,Che12}. For $E_\\text{sym}(\\rho_0)$ and $L $, all existing constraints extracted so far from terrestrial laboratory measurements and astrophysical observations are found to be essentially consistent with $E_{\\text{sym}}({\\rho _{0}}) = 32 \\pm 3$ MeV and $L = 50 \\pm 20$ MeV~\\cite{Che12a,Li12}. While these lower-order bulk characteristic parameters have been relatively well determined or in significant progress, our knowledge on the higher-order bulk characteristic parameters remains very limited. Following $E_0(\\rho_0)$, $K_0$, $E_\\text{sym}(\\rho_0)$ and $L $, the next bulk characteristic parameter should be the skewness coefficient $J_0$ (also denoted as $K'$ or $Q_0$ in some literature) of SNM, which is related to the third-order density derivative of the binding energy per nucleon of SNM at $\\rho_0 $. The higher-order bulk characteristic parameter $J_0$ is expected to be important for the high density behaviors of nuclear matter EOS and thus may play an essential role in heavy ion collisions (HIC), the structure and evolution of neutron stars, supernova explosion, and gravitational wave radiation from merging of compact stars. To our best knowledge, so far there is very little experimental information on the $J_0$ parameter, and it is thus of great interest and critical importance to constrain the $J_0$ parameter. Within the nonlinear relativistic mean field (RMF) model, we demonstrate in this work that combining the experimental constraints on the pressure of SNM at supra-saturation densities from flow data in HIC and the astrophysical observation of a large mass neutron star PSR J0348+0432, one can obtain a strong constraint on the $J_0$ parameter. ",
        "conclusions": "For the Lagrangian in Eq.~(\\ref{NLRMF}), the properties of infinite nuclear matter is uniquely determined by $f_{\\sigma}=g_{\\sigma}/m_{\\sigma}$, $b_{\\sigma}$, $c_{\\sigma}$, $f_{\\omega}=g_{\\omega}/m_{\\omega}$, $c_{\\omega}$, $f_{\\rho}=g_{\\rho}/m_{\\rho}$, $\\Lambda_{\\mathrm{V}}$, and $M$. If the nucleon mass in vacuum is set to be $M = 939$ MeV, one then has totally seven parameters to determine the properties of infinite nuclear matter in the nonlinear RMF model. Following the correlation analysis method proposed in Ref.~\\cite{Che10} within the Skyrme-Hartree-Fock (SHF) approach, instead of using directly the seven microscopic parameters, i.e., $f_{\\sigma}$, $b_{\\sigma}$, $c_{\\sigma}$, $f_{\\omega}$, $c_{\\omega}$, $f_{\\rho}$, $\\Lambda_{\\mathrm{V}}$, one can determine their values explicitly in terms of seven macroscopic quantities, i.e., $\\rho _{0}$, $E_{0}(\\rho_{0})$, $K_{0}$, $J_{0}$, $M_{\\mathrm{dirac}}^{\\ast 0}\\equiv M_{\\mathrm{dirac}}^{\\ast}(\\rho_{0})$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$, and $L(\\rho_{\\mathrm{c}})$ where $\\rho_{\\mathrm{c}}$ is the cross density whose value is fixed in this work at $0.11\\,\\mathrm{fm}^{-3}$~\\cite{Zha13}. Then, by varying individually these macroscopic quantities within their known ranges, one can examine more transparently the correlation of nuclear matter properties with each individual macroscopic quantity. Recently, this correlation analysis method has been successfully applied to study the neutron skin~\\cite{Che10,Zha13} and GMR of finite nuclei~\\cite{Che12}, the higher-order bulk characteristic parameters of ANM~\\cite{Che11a}, and the relationship between the nuclear matter symmetry energy and the symmetry energy coefficient in the mass formula~\\cite{Che11}.  \\renewcommand*\\figurename{\\small Figure} \\begin{figure}[tbh] \\centering \\includegraphics[scale=0.80]{PressCorr.eps} \\caption{Pressure of SNM at $\\rho=4\\rho_0$ from the nonlinear RMF model with the FSUGold interaction by varying individually $\\rho_{0}$ (a), $E_{0}(\\rho_{0})$ (b), $M_{\\mathrm{dirac}}^{\\ast 0}$ (c), $K_{0}$ (d), and $J_{0}$ (e).} \\label{PressCorr} \\end{figure}  To examine the correlation of pressure of SNM at supra-saturation densities with each macroscopic quantity, we show in Fig.~\\ref{PressCorr} the pressure of SNM $P(\\rho)$ at $\\rho=4\\rho_0$ from the nonlinear RMF model with the FSUGold interaction~\\cite{Tod05} by varying individually $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, and $J_{0}$ within their empirical uncertain ranges, namely, varying one quantity at a time while keeping all others at their default values in FSUGold for which we have $\\rho _{0}=0.148$ fm$^{-3}$, $E_{0}(\\rho_{0})=-16.3$ MeV, $M_{\\mathrm{dirac}}^{\\ast 0}=0.61M$, $K_{0}=230$ MeV, $J_{0}=-522.6$ MeV, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})=27.11$ MeV, and $L(\\rho_{\\mathrm{c}})=49.97$ MeV. It should be mentioned that the pressure of SNM is independent of the values of $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ and $L(\\rho_{\\mathrm{c}})$. It is seen from Fig.~\\ref{PressCorr} that the pressure of SNM $P(\\rho)$ at $\\rho=4\\rho_0$ increases with $\\rho_{0}$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$ and $J_{0}$ while decreases with $E_{0}(\\rho_{0})$. In particular, the pressure of SNM $P(\\rho)$ at $\\rho=4\\rho_0$ displays a specially strong correlation with $J_{0}$. We note that the pressure of SNM $P(\\rho)$ at other supra-saturation densities exhibits similar correlations with $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, and $J_{0}$. These features indicate that the pressure of SNM at supra-saturation densities is sensitive to the $J_0$ value, and thus the experimental constraints on the pressure of SNM at supra-saturation densities may provide important information on the $J_0$ value.  Since the pressure of SNM at supra-saturation densities is sensitive to the $J_0$ value, the maximum mass $M_{\\text{max}}$ of static neutron stars is also expected to be sensitive to the $J_0$ value. The mass and radius of static neutron stars can be obtained from solving the Tolman-Oppenheimer-Volkoff (TOV) equations with a given neutron star matter EOS. A neutron star generally contains core, inner crust and outer crust from the center to surface. In this work, for the core where the baryon density is larger than the core-curst transition density $\\rho_{\\text{t}}$, we use the EOS of $\\beta$-stable and charge neutral npe$\\mu$ matter obtained from the nonlinear RMF model. In the inner crust with densities between $\\rho_{\\text{out}}$ and $\\rho _{\\text{t}}$ where the nuclear pastas may exist,  we construct its EOS (pressure $P$ as a function of energy density $\\mathcal{E}$) according to $P=a+b\\mathcal{E} ^{4/3}$ because of our poor knowledge about its EOS from first principle~\\cite{Hor03,XuJ09}. The $\\rho_{\\text{out}}=2.46\\times 10^{-4}$ fm$^{-3}$ is the density separating the inner from the outer crust. The constants $a$ and $b$ are then easily determined by the pressure and energy density at $\\rho _{\\text{t}}$ and $\\rho _{\\text{out}}$~\\cite{XuJ09}. In this work, the $\\rho _{\\text{t}}$ is determined self-consistently within the nonlinear RMF model using the thermodynamical method (See, e.g., Ref.~\\cite{Cai12} for the details). In the outer crust with $6.93\\times 10^{-13}$ fm$^{-3}<\\rho <\\rho _{\\text{out}}$, we use the EOS of BPS \\cite{BPS71,Iida97}, and in the region of $4.73\\times 10^{-15}$ fm$^{-3}<\\rho <$$6.93\\times 10^{-13}$ fm$^{-3}$ we use the EOS of FMT~\\cite{BPS71}.  Similarly as in Fig.~\\ref{PressCorr}, we plot in Fig.~\\ref{MassCorr} the maximum mass $M_{\\text{max}}$ of static neutron stars from the nonlinear RMF model with the FSUGold interaction by varying individually $\\rho _{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, $J_{0}$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$, and $L(\\rho_{\\mathrm{c}})$ within their empirical uncertain ranges. Indeed, one can see that the $M_{\\text{max}}$ displays a very strong positive correlation with the $J_0$ parameter. In addition, the $M_{\\text{max}}$ exhibits weak positive correlation with the $M_{\\mathrm{dirac}}^{\\ast 0}$ and $K_{0}$, and weak negative correlation with the $\\rho _{0}$ and $E_{0}(\\rho_{0})$. It is interesting to see that the $M_{\\text{max}}$ is essentially independent of the values of $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ and $L(\\rho_{\\mathrm{c}})$, implying that, in the nonlinear RMF model, the $M_{\\text{max}}$ is basically determined by the isoscalar part of the nuclear matter EOS. These features indicate that the observed largest mass of neutron stars may put important constraint on the $J_0$ value.  \\begin{figure}[tbh] \\includegraphics[scale=0.81]{MassCorr.eps} \\caption{Maximum mass of static neutron stars from the nonlinear RMF model with the FSUGold interaction by varying individually $\\rho _{0}$ (a), $E_{0}(\\rho_{0})$ (b), $M_{\\mathrm{dirac}}^{\\ast 0}$ (c),  $K_{0}$ (d), $J_{0}$ (e), $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ (f), and $L(\\rho_{\\mathrm{c}})$ (g).} \\label{MassCorr} \\end{figure}  Experimentally, the pressure of SNM at supra-saturation densities (from $2\\rho_0$ to about $5\\rho_0$) has been constrained by measurements of collective flows in HIC~\\cite{Dan02}, which is shown as a band in the left window of Fig.~\\ref{FlowNStar}. In the nonlinear RMF model, if one only changes the $J_0$ value while the other $6$ macroscopic quantities are kept at their values in the FSUGold interaction, one can find that the $J_0$ should be in the range of $-985 \\mathrm{MeV}\\leq J_0\\leq -327 \\mathrm{MeV}$ to be consistent with the flow data in HIC~\\cite{Dan02}. However, keeping the other $6$ macroscopic quantities at their values in the FSUGold interaction is obviously a strong assumption because the extraction of the $J_0$ value from the flow data in HIC will also depends on the values of $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$, which can be varied within their empirical uncertain ranges. For the nonlinear RMF model, we use in this work the following empirical uncertain ranges for these macroscopic quantities, i.e., $\\rho_0=0.153\\pm0.008\\,\\mathrm{fm}^{-3}$, $E_0(\\rho_0)=-16.2\\pm0.3\\,\\mathrm{MeV}$, $M_{\\mathrm{dirac}}^{\\ast 0}/M=0.61\\pm0.04$, and $K_0=230\\pm20\\,\\mathrm{MeV}$, which represent the typical uncertain ranges known or predicted from different interactions in the nonlinear RMF model~\\cite{Che07}.  Based on the pressure of SNM constrained by flow data in HIC~\\cite{Dan02}, to extract the upper limit of the $J_0$ value, one should use the values of $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$ that make the resulting pressure of SNM as small as possible when $J_0$ is fixed. This can be obtained by using $\\rho_0=0.145\\,\\mathrm{fm}^{-3}$, $E_0(\\rho_0)=-15.9\\,\\mathrm{MeV}$, ${M_{\\rm{dirac}}^{\\ast 0}}/{M}=0.57$, and $K_0=210\\,\\mathrm{MeV}$, denoted as set ``S\", since the pressure of SNM $P(\\rho / \\rho_0)$ at supra-saturation densities increases with $\\rho_0$, $M_{\\mathrm{dirac}}^{\\ast 0}$ and $K_{0}$ while decreases with $E_{0}(\\rho_{0})$ as shown in Fig.~\\ref{PressCorr}. With the set ``S\" for $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$, one can find the upper limit of $J_0= -10$ MeV, which is indicated by solid line in the left window of Fig.~\\ref{FlowNStar}. For $J_0 > -10$ MeV, the model would over-predict the pressure of SNM constrained by flow data in HIC~\\cite{Dan02}. Similarly, one can obtain the lower limit of $J_0$ by using the values of $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$ that make the resulting pressure of SNM as large as possible when $J_0$ is fixed, and this can be obtained with $\\rho_0=0.161\\,\\mathrm{fm}^{-3}$, $E_0(\\rho_0)=-16.5\\,\\mathrm{MeV}$, ${M_{\\rm{dirac}}^{\\ast 0}}/{M}=0.65$, and $K_0=250\\,\\mathrm{MeV}$, denoted as set ``H\". Using the set ``H\" for $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$, one can extract the lower limit of $J_0= -1280$ MeV, which is indicated by dashed line in the left window of Fig.~\\ref{FlowNStar}. The model would under-predict the pressure of SNM constrained by flow data in HIC~\\cite{Dan02} if $J_0 < -1280$ MeV. Therefore, from the pressure of SNM constrained by flow data in HIC~\\cite{Dan02}, one can extract the constraint of $-1280 \\text{ MeV} \\le J_0 \\le -10 \\text{ MeV}$.  \\renewcommand*\\figurename{\\small Figure} \\begin{figure}[tbh] \\centering \\includegraphics[width=8.5cm]{FlowNStar.eps} \\caption{(Color online) Left window: Pressure of SNM as a function of baryon density. The solid (dashed) line is the prediction from the nonlinear RMF model with $J_0=-10$ ($-1280$) MeV and the set ``S (H)\" for $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, and $K_{0}$. The band represents the constraints from flow data in HIC~\\cite{Dan02}. Right window: The maximum mass of static neutron stars as a function of $J_0$ in the nonlinear RMF model with the set ``NS-H\" for $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ and $L(\\rho_{\\mathrm{c}})$. The band represents mass $2.01\\pm 0.04M_{\\odot }$ for PSR J0348+0432~\\citep{Ant13}.} \\label{FlowNStar} \\end{figure}  Recently, a new neutron star PSR J0348+0432 with a mass of $2.01\\pm0.04M_{\\odot}$ was discovered~\\cite{Ant13}, and this neutron star is only the second pulsar with a precisely determined mass around $2M_{\\odot}$ after PSR J1614-2230~\\citep{Dem10} and sets a new record of the maximum mass of neutron stars. The lower mass limit of $1.97M_{\\odot}$ for PSR J0348+0432 thus may set a lower limit of $J_0$ below which the model cannot predict a neutron star with mass equal or above $1.97M_{\\odot}$. To extract the lower limit of $J_0$ from the observed heaviest neutron star PSR J0348+0432, one can use the values of $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ and $L(\\rho_{\\mathrm{c}})$ that make the resulting maximum mass of neutron stars as large as possible when $J_0$ is fixed, and from Fig.~\\ref{MassCorr} this can be obtained with $\\rho_0=0.145\\,\\mathrm{fm}^{-3}$, $E_0(\\rho_0)=-16.5\\,\\mathrm{MeV}$, $K_0=250\\,\\mathrm{MeV}$, ${M_{\\rm{dirac}}^{\\ast 0}}/{M}=0.65$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})=26\\,\\mathrm{MeV}$ and $L(\\rho_{\\mathrm{c}})=60\\,\\mathrm{MeV}$, denoted as set ``NS-H\". This leads to a lower limit of $J_0= -494$ MeV as shown in the right window of Fig.~\\ref{FlowNStar} where the maximum mass of neutron stars is plotted as a function of $J_0 $ when the other $6$ macroscopic quantities are fixed at their values as in set ``NS-H\". For $J_0 < -494$ MeV, the maximum mass of static neutron stars predicted in the nonlinear RMF model would be always smaller than $1.97M_{\\odot}$. It should be pointed out that here the interior of neutron stars has been assumed to be npe$\\mu$ matter. New degrees of freedom such as hyperons or/and quark matter that could be present in the interior of neutron stars usually soften the EOS of neutron star matter and thus a larger $J_0$ value would be necessary to obtain a neutron star with mass of $1.97M_{\\odot}$. Therefore, including the new degrees of freedom in neutron stars will be consistent with the constraint of $J_0 \\ge -494$ MeV.  Combining the constraint of $-1280 \\text{ MeV} \\le J_0 \\le -10 \\text{ MeV}$ from the pressure of SNM constrained by flow data in HIC~\\cite{Dan02} which favors a smaller $J_0$ value and the constraint of $J_0 \\ge -494$ MeV from the recently discovered heaviest neutron star PSR J0348+0432~\\citep{Ant13} which favors a larger $J_0$ value, one can extract the following constraint for the $J_0$ parameter \\begin{equation}\\label{cons_J0} -494\\,\\mathrm{MeV}\\leq J_0\\leq -10\\,\\mathrm{MeV}. \\end{equation} This constraint represents a conservative extraction based on flow data in HIC~\\cite{Dan02} and the recently discovered heaviest neutron star PSR J0348+0432~\\citep{Ant13} in the nonlinear RMF model since we have not considered the possible correlations existed among the $\\rho_{0}$, $E_{0}(\\rho_{0})$, $M_{\\mathrm{dirac}}^{\\ast 0}$, $K_{0}$, $E_{\\rm{sym}}(\\rho_{\\mathrm{c}})$ and $L(\\rho_{\\mathrm{c}})$ and set simultaneously their values in the boundary of their empirical uncertain ranges. Considering the possible correlations will further narrow the above constraint. The above constraint indicates that, if the $J_0$ value is out of the region $-494\\,\\mathrm{MeV}\\leq J_0\\leq -10\\,\\mathrm{MeV}$, the nonlinear RMF model either cannot predict the pressure of SNM constrained by flow data in HIC~\\cite{Dan02} or cannot describe the recently discovered heaviest neutron star PSR J0348+0432~\\citep{Ant13}.  \\renewcommand*\\figurename{\\small Figure} \\begin{figure}[tbh] \\centering \\includegraphics[width=7.0cm]{J0.eps} \\caption{(Color online) Comparison between the constraint of $J_0$ extracted in the present work and those from the existing literature: Farine {\\it et al.}~\\cite{Far97}, Steiner {\\it et al.}~\\cite{Ste10}, Chen~\\cite{Che11a}, and Meixner {\\it et al.}~\\cite{Mei13}.} \\label{J0} \\end{figure}  In Fig.~\\ref{J0}, we show the comparison of the constraint of $J_0$ found in our analysis with those found with other analysis and/or other methods~\\cite{Far97,Ste10,Che11a,Mei13}, including the constraint of $J_0 = -700 \\pm 500$ MeV obtained by Farine {\\it et al.}~\\cite{Far97} from the analysis of nuclear GMR, the constraint of $J_0 = -280^{+72}_{-410}$ ($-500^{+170}_{-290}$) MeV obtained by Steiner {\\it et al.}~\\cite{Ste10} from analyzing a heterogeneous data set of six neutron stars using a Markov chain Monte Carlo algorithm within a Bayesian framework by assuming $r_{\\text{ph}} \\gg R$ ($r_{\\text{ph}}= R$) where $r_{\\text{ph}}$ is the photospheric radius at the time the flux is evaluated and $R$ is the stellar radius, the constraint of $J_0 = -355 \\pm 95$ MeV deduced by Chen~\\cite{Che11a} based on a correlation analysis method within SHF energy density functional, and the constraint of $J_0 = -390 \\pm 90$ MeV deduced by Meixner {\\it et al.}~\\cite{Mei13} who used the similar method as Chen~\\cite{Che11a}. It is seen that the constrained region of $J_0$ obtained in the present work has a remarkable overlap with those existing in the literature. In particular, our present constraint from the relativistic model is nicely consistent with the constraints deduced from the non-relativistic SHF approach and they all indicates that the $J_0$ parameter should be larger than about $-500$ MeV."
    },
    "1402/1402.3979_arXiv.txt": {
        "abstract": "Bajc-Melfo(\\textbf{BM})   two field  ($S,\\phi$)  superpotentials define metastable F-term supersymmetry breaking   vacua suitable as    hidden sectors    for calculable and realistic family and Grand Unification  unification models. The undetermined vev $<S_s>$ of the Polonyi field that  breaks Supersymmetry  can be fixed either by coupling to N=1 Supergravity or by radiative corrections. \\textbf{BM}   hidden sectors extend   to symmetric multiplets $(S,\\phi)_{ab}$ of a gauged  $O(N_g)$ family symmetry, broken at the GUT scale,  so that the $O(N_g)$ charged component vevs   $<\\hat S_{ab}>$ are  also undetermined    before  accounting for the $O(N_g)$ D-terms: which fix them by  cancellation against     D-term contributions from the visible sector.  This  facilitates Yukawon Ultra Minimal GUTs(YUMGUTs) proposed in [C.S.Aulakh and C.K.Khosa, Phys.Rev.D 90,045008(2014)] by   relieving the visible sector from the need to give null  D-terms for the family symmetry  $ O(N_g)$. We analyze  symmetry breaking and and spectra of the hidden sector fields in the Supergravity resolved case  when  $N_g=1,2,3$. Besides the Polonyi field $S_s$, most of the superfields $\\hat S_{ab}$ remain light, with fermions getting masses only from loop corrections. Such  modes  may yield novel dark matter lighter than 100 GeV.  Possible  Polonyi and moduli problems associated with the  the fields  $S_{ab}$ call  for   detailed investigation of loop effects due to the Yukawa and gauge interactions in the hidden sector and of post-inflationary field relaxation dynamics. ",
        "introduction": "Supersymmetry imposes remarkable restrictions on its own spontaneous violation and thus   severely constrains   models of BSM physics. Traditionally\\cite{cremmer,ohta,nath,weinberg82,halllykkwein} gravity mediation has been used to generate phenomenologically acceptable soft Susy breaking terms for globally supersymmetric GUTs. This requires supersymmetry breaking in a hidden sector with no superpotential couplings to the `visible' sector in which we live. The GUT Minkowski space vacuum is metastable\\cite{weinberg82} in the sense of there being a Susy preserving vacuum with lower (negative) energy, but it is thought that it cannot decay to the lower vacuum by tunnelling; or at any rate the timescale for doing so is much larger than the age of the universe.  The vacuum in the gravity mediated scenario is defined as the perturbation of a Global Susy preserving minimum by a hidden sector with a \\emph{global} minimum which breaks supersymmetry spontaneously due to specialized superpotentials(Polonyi,O'Raifeartaigh \\cite{polo,raff} etc) or due to  strong dynamics(gaugino condensation etc). After long struggles to tame the hard problem of finding phenomenologically usable  Susy breaking global minima it was re-emphasized\\cite{dimodvalrattgiu,intrillshih} that metastable i.e local minima could serve just as well as global minima, since the unifying theory would anyway  have vacua at lower energy\\cite{weinberg82} than the phenomenologically relevant Minkowski vacuum on which the MSSM lived.   This realization was liberating for phenomenologists because  models defined on supersymmetry breaking vacua which are  only local minima offer  a much larger range of candidate vacua. Various strongly coupled supersymmetric gauge theories \\cite{intrillshih,Lutyagashe}were studied to provide the Susy breaking dynamics albeit now considering also metastable vacua. Unfortunately even the metastable dynamical Susy breaking models seem  so arbitrarily elaborate as to rival the complexity of the phenomenological theory (GUT,MSSM etc) for which they were supposedly to provide the service of Supersymmetry breaking: making testability of both a moot issue.   Recently Bajc and Melfo (BM)\\cite{BM} suggested that,  from a `calculable unification'  viewpoint, it would be more interesting to consider Susy breaking metastable vacua of simple two field hidden sector superpotentials. BM models  are related to the classic O'Raifeartaigh susy breaking  superpotentials and are variants of a class of models with   susy breaking local minima studied systematically earlier  in \\cite{ray}. Such superpotentials typically have flat directions and there is a longstanding  vision whereby the determination of such flat directions(running out of susy breaking minima) by radiative corrections\\cite{wittenhilo} or supergravity\\cite{ovrab,csahilo} might profitably fix the undetermined vacuum expectation values(VEVs) at a high Unification scale. Bajc and Melfo took up this challenge and formulated models where the  undetermined VEVs were those of \\textbf{24}  or \\textbf{75} irreps of an $SU(5)$ unified theory and the flat directions were determined by radiative corrections. The use of of N=1 Supergravity potentials to lift the flat directions rolling out of the Bajc-Melfo   Susy breaking local minimum has  not   been  analyzed.  From a Grand Unified viewpoint the only credible sign of flavour unification so far so far is  third family  Yukawa unification at large $\\tan \\beta$ in SO(10)   at  scales of order $M_X > 10^{16}$ GeV. On the other hand   minimal Supersymmetric SO(10)  GUTs\\cite{aulmoh,ckn,abmsv,bmsv,ag2, gmblm,blmdm,nmsgut}  have been shown to be completely realistic and even to suppress $d=5$ proton decay\\cite{gutupend,bstabhedge} by a novel and generic mechanism based on the necessary  emergence of a pair of light doublet Higgs in the effective MSSM. Motivated by these two clues  we exploited  the defining(i.e  \\boot ) representation of MSGUTs(which gets large vevs yet couples to matter fermions)  to design a \\emph{renormalizable}  Grand Unified Yukawon/familion model\\cite{yukawon} in which we generate hierarchical matter fermion  Yukawas from flavour bland  parameters of the MSGUT extended by an $O(N_g)$ family symmetry. Family gauge symmetry is  broken at the unification scale and the visible sector GUT Higgs fields  effloresce  into symmetric ($\\mathbf{210,126,\\oot,10}$)  or anti-symmetric ($\\mathbf{120}$ ) irreps of the family group (as dictated by the properties of bilinears in the  \\textbf{16}-plet matter  irreps in SO(10) GUTs), while the matter \\textbf{16-}plet Yukawa couplings become flavour bland.  The neat and exactly soluble\\cite{abmsv} gauge symmetry breaking ($SO(10)\\rightarrow \\rm{MSSM}$) of the family singlet MSGUT Higgs system then generalizes straightforwardly for the family symmetry variant MSGUT Higgs multiplets. The supersymmetric   breaking of family symmetry will be discussed in detail in this paper. Extraction of the light MSSM Higgs from the plethora of MSSM doublets present in the $O(N_g)\\times SO(10)$ GUT  via calculation of the `Higgs fractions'\\cite{abmsv,ag1,ag2,nmsgut} (which specify the make up of the MSSM Higgs) determined by flavour blind GUT parameters generates candidate  flavour structures.  In implementing this flavour generation scenario an obstacle arises. Although the solution of the F and D term conditions for the GUT multiplets in the SO(10) sector with flavour  non-diagonal VEVs is straightforward, it is very hard or impossible to ensure that the same configurations simultaneously cancel the family symmetry D-terms. It thus becomes necessary to introduce additional fields to soak up the the contribution   of the GUT VEVs   to the $O(N_g)$ D terms. Since minimality, simplicity and solubility of the GUT Higgs sector symmetry breaking is a central virtue of MSGUTs   it is natural to ask whether the  additional fields might not be located in a hidden sector with its own flat directions arising from supersymmetry breaking in a way which did not interfere with the already accomplished MSGUT symmetry breaking\\cite{aulmoh,ckn,abmsv}. Besides making a completely unexpected connection between family symmetry  and supersymmetry breaking this idea is straightforward to implement in a context where the BM flat directions are lifted by a combination of supergravity and $O(N_g)$ gauge D term minimization effects.  A notable phenomenological implication of our model  is the presence of  weakly coupled light fields(`moduli'). In particular, in the   supergravity mediated scenario presented here there is a Polonyi (scalar) field (Planck scale vev, weak scale mass and intermediate scale susy breaking F-term). In addition there are other light fermions and scalars  associated with the flat directions of the family non singlet BM fields. These modes may be interesting light( $<50$ GeV) Susy WIMP dark matter  of the sort perhaps indicated by the DAMA/LIBRA experiment \\cite{damalibra} and so far completely missing from MSGUT spectra. On the other hand this clutch of modes may also  pose the typical problems associated with `moduli' fields found in other fundamental theories\\cite{modulicon}. Recall that the Polonyi problem arises  due to late decay of the oscillations of the supersymmetry breaking  Polonyi field while it settles into the zero temperature potential's    minimum because of the suppression of   its couplings by the Planck scale. The dark matter candidates' role  in cosmogony also needs to be checked for consistency. In view of the new Yukawa and gauge interactions of the hidden sector the cosmological implications require a separate detailed study and thus  should not prematurely prejudice the novel flavour generation aspects of our model. In this paper we give the relevant details for the symmetry breaking, VEVs and masses in the BM sector  for direct use in our `yukawonification' \\cite{yukawon} models. We will return to a study of cosmological constraints and alternative models elsewhere. ",
        "conclusions": "In this paper we have shown  how, gravity mediation of Bajc-Melfo type calculable   Supersymmetry breaking at metastable vacua in a hidden sector containing fields variant under  a family gauge symmetry, ensures the viability of  GUT and family symmetry breaking in the  visible sector of a ``Grand Yukawonification ''\\cite{yukawon} model. In such  models one aims at generating the observed hierarchical fermion  flavour structure from a gauged family symmetry  model with only generation blind couplings. The special role of the BM supersymmetry breaking is that it provides flat directions in both the   $O(N_g)$ singlet and the gauge variant parts of a symmetric chiral supermultiplet $S_{ab}=S_{ba}$.   Since it is very difficult to make the contribution of the visible sector GUT fields to the $O(N_g)$ D-terms vanish,   the $\\hat S$ flat direction performs the invaluable function of cancelling this contribution \\emph{without disturbing the symmetry breaking in the visible sector}. The determination of the viability of the ``Grand Yukawonification'' model then becomes a matter of searching the relatively  small remaining parameter space for viable parameter sets that fit the fermion data at $M_X$ while taking account of threshold corrections at low and high scales and while respecting constraints on crucial quantities like the proton lifetime\\cite{ag2,nmsgut,bstabhedge}. Note that in this approach not only are the hard parameters of the visible sector superpotential reduced by replacement of the flavoured parameters by bland family symmetric ones  but also the soft supersymmetry breaking parameters are determined by the two parameters of the hidden sector superpotential and the Planck scale.  The structure used entails yet further stringent constraints since the masslessness of the   moduli multiplets  $\\hat{S}_{ab}$ before supersymmetry breaking  implies  the existence of $N_g(N_g+1)/2 -1$ SM singlet fermions generically lighter than the gravitino mass scale and possibly as light as a few GeV. In addition the Polonyi mode $S_s$ may  also lead  to difficulties in the cosmological scenario. Thus such modes can be both a boon and a curse for familion GUT models. A boon because generic Susy GUT models are hard put if asked to provide Susy WIMPs of mass   below 100 GeV as CDM candidates as suggested by the DAMA/LIBRA experiment\\cite{damalibra}. A curse because there are strong constraints on the existence of such light moduli which normally demand that their mass be rather large ($> 10$ TeV) due to the robust   cosmological(`Polonyi') problems arising from   decoupled modes with Planck scale VEVs\\cite{modulicon}. In contrast to the simple Polonyi model and String moduli, the BM moduli  have explicit   couplings  to light fields through family D-term mixing and loops. Moreover the MSGUT scenario favours\\cite{nmsgut,gutupend,bstabhedge}  large gravitino masses $> 5-50 $ TeV. Thus  the Polonyi  and moduli problems may be evaded. In any case the cosmology need be considered seriously only after we have shown that the MSSM fermion spectrum is indeed generated by the \"Yukawonified\" NMSGUT\\cite{yukawon}."
    },
    "1402/1402.1184_arXiv.txt": {
        "abstract": "We present far-infrared and sub-millimeter images of the \\ec\\ debris disk system obtained with {\\it Herschel} and SCUBA-2, as well as {\\it Hubble Space Telescope} visible and near-infrared coronagraphic images. In the 70\\,\\mic\\ {\\it Herschel} image, we clearly separate the thermal emission from the warm and cold belts in the system, find no evidence for a putative dust population located between them, and precisely determine the geometry of the outer belt. We also find marginal evidence for azimuthal asymmetries and a global offset of the outer debris ring relative to the central star. Finally, we place stringent upper limits on the scattered light surface brightness of the outer ring. Using radiative transfer modeling, we find that it is impossible to account for all observed properties of the system under the assumption that both rings contain dust populations with the same properties. While the outer belt is in reasonable agreement with the expectations of steady-state collisional cascade models, albeit with a minimum grain size that is four times larger than the blow-out size, the inner belt appears to contain copious amounts of small dust grains, possibly below the blow-out size. This suggests that the inner belt cannot result from a simple transport of grains from the outer belt and rather supports a more violent phenomenon as its origin. We also find that the emission from the inner belt has not declined over three decades, a much longer timescale than its dynamical timescale, which indicates that the belt is efficiently replenished. ",
        "introduction": "Questions regarding the origin of the Solar System have pervaded the history of astronomy. Thorough studies of the architecture and dynamical history of our own Solar System \\citep[e.g.,][]{horner14} and the discovery of thousands of extrasolar planets \\citep{howard13} now provide an opportunity to test theoretical models and to piece together the story of planetary systems. Another natural laboratory of this history consists of debris disks, optically thin circumstellar disks composed mainly of dust surrounding Main Sequence stars \\citep{zuckerman01}. In these systems, micron-sized dust grains, the only ones that can be probed directly, are continually blown out by radiation forces and replenished primarily via the grinding of planetesimals. Debris disks are observed at a rate of 15-20\\% around Main Sequence stars, with increased frequency around younger stars \\citep[][; Thureau et al., in prep.; Sibthorpe et al., in prep.]{su06, trilling08, carpenter09, eiroa13}. Therefore, they represent a crucial window into the planet formation process, each system providing a snapshot at a particular age and a chance to understand its dependency on other stellar parameters such as mass and metallicity. Understanding the diversity of debris disk systems is an important step towards placing our own Solar System in a global context and drawing a complete picture of planet formation.  Most debris disk systems are consistent with a steady-state collisional cascade within a gradually decaying reservoir of planetesimals, the canonical theoretical framework for these objects \\citep{wyatt08}. However, this scenario is seriously challenged by systems showing very strong emission from hot/warm (i.e., close in) dust, especially if the system is relatively old \\citep[$\\gtrsim 100$\\,Myr, i.e. past the epoch of terrestrial planet formation; for a recent review see][]{chambers11}. Instead, such systems are most likely transient remnants of a recent or ongoing event, either a rare catastrophic collision between very large planetesimals, a Late Heavy Bombardment-like phenomenon or a continuous replenishment via comets from a massive reservoir located much further out \\citep{lisse12, kennedy13}. Dust belts located in the planet-forming region, i.e., within a few AU of the central star, can therefore arise from several paths and only detailed studies of these systems can inform on the likelihood of each scenario. Of particular interest is whether these particularly dusty warm debris disks are associated with a colder dust belt that could serve to feed dust grains to the innermost regions of the systems. Careful analysis of {\\it Spitzer} mid-infrared spectroscopy has revealed that debris systems containing two separate dust belts, akin to the situation in our Solar System albeit with much more copious amounts of dust, are very common \\citep[][Chen et al., submitted]{morales11}. The large physical separation between the two dust belts in these systems is suggestive of the presence of planetary bodies between them \\citep{su13}, although it could also be explained by dust produced from planetesimals on initially high eccentricity orbits \\citep{wyatt10}.  Debris disks are usually discovered by the infrared excess in the spectral energy distributions (SEDs) caused by the thermal emission of the dust. However, interpretations of the SED are fraught with an inherent degeneracy, in that a population of larger grains closer to the central star will exhibit the same total luminosity and equilibrium temperature as a population of smaller grains located further out. Indeed, spatially-resolved images of debris disks frequently reveal larger radii than would be inferred from pure blackbody emission \\citep{rodriguez12, booth13, morales13}. Mid-infrared spectroscopy revealing silicate emission features can help break this degeneracy but is only available for a very limited number of targets \\citep[e.g.,][]{chen06}. Thus, resolved imagery is essential to determine the dust properties, since it provides critical constraints on the location of the grains. Furthermore, spatially resolved images of debris disks can provide indirect evidence for the presence of planet-mass bodies that can interact gravitationally and shepherd the small dust grains \\cite[e.g.,][]{kalas08, lagrange09}.  \\ec\\ is an F2, 1.4\\,$M_\\odot$, Main Sequence star located only 18.3\\,pc away \\citep{vanleeuwen07}, whose age has been estimated to be 1--2\\,Gyr \\citep{ibukiyama02, mallik03, vican12}. Significant infrared excess emission was first detected with {\\it IRAS} \\citep{stencel91}, and many subsequent studies over a broad wavelength range have targeted and analyzed this remarkable system. Most intriguingly, the SED of \\ec\\ exhibits a clear two-peaked excess, which is most easily explained by the fact that the circumstellar dust is distributed in two well separated belts. Indeed, \\cite{wyatt05} presented sub-millimeter maps that resolved the system for the first time and showed that the cold dust belt, located $\\sim$150\\,AU from the central star, could not be responsible for the copious amounts of mid-infrared emission \\citep[see also][]{beichman06, rhee07}. Instead, a second belt, located within $\\lesssim$3\\,AU of the host star \\citep{smith08, smith09}, is the source of this emission at shorter wavelengths. Detailed analyses of the dust composition in this inner belt have been enabled by the spectroscopic capabilities of {\\it Spitzer}, showing a rich mineralogy and the unambiguous presence of grains as small as 1\\,\\mic\\ \\citep{chen06, lisse12}. Most importantly, the mid-infrared excess produced by this dust belt is remarkably large considering the old age of the system, indicating that the gradual erosion via collisional grinding scenario is untenable \\citep[e.g.,][]{wyatt07}. Throughout this paper, we interchangeably refer to the inner (outer) as the warm (cold) ring/belt.  As the 63$^{\\rm rd}$ closest F-type star to the Sun \\citep{phillips10}, \\ec\\ was included in the volume-limited Disc Emission via a Bias-free Reconnaissance in the Infrared/Sumillimetre survey (DEBRIS, Matthews et al., in prep.) which uses far-infrared continuum imaging with {\\it Herschel} \\citep{pilbratt10} to obtain more sensitive and higher resolution images of debris disk systems than previous observations in this domain. Early ``Science Demonstration\" 100 and 160\\,\\mic\\ images of \\ec\\ clearly resolved the outer dust belt \\citep{matthews10} and warranted follow-up observations to take full advantage of the capabilities of {\\it Herschel}. In this paper, we present new {\\it Herschel} images of \\ec\\ at wavelengths ranging from 70 to 500\\,\\mic, as well as a new ground-based 850\\,\\mic\\ map of the system obtained as part of the SCUBA-2 Observations of Nearby Stars survey \\citep[SONS,][]{panic13}. The 70\\,\\mic\\ image presented here is the highest resolution image of the system to date and it allows us to clearly disentangle the emission from both dust belts. We are thus able to determine precisely the geometry of the cold outer belt and to constrain the far-infrared emission properties of the warm inner belt. In turn, this enables us to constrain the dust properties in the outer belt and to compare them to those of the inner belt, providing crucial insight on the nature and current state of the system. We also present new visible and near-infrared {\\it Hubble Space Telescope} ({\\it HST}) images that allow us to place an upper limit on the amount of starlight that scatters off the outer belt.  The organization of this paper is as follows. In Section\\,\\ref{sec:obs} we present the new observations of \\ec\\ and the data reduction methods, as well as a re-analysis of its mid-infrared {\\it Spitzer} spectrum. In Section\\,\\ref{sec:results}, we present the images as well as some basic analysis of both the SED, using modified blackbody models, and the images, using simple geometric models. In Section\\,\\ref{sec:model}, we show the results of a full radiative transfer model assuming a warm inner and cold outer belt, from which the dust properties in both dust belts are constrained, and we discuss the implications of our results in Section\\,\\ref{sec:discus}. ",
        "conclusions": "As part of the DEBRIS survey, we have obtained new {\\it Herschel} 70--500\\,\\mic\\ images of the \\ec\\ debris disk system, as well as a new ground-based 850\\,\\mic\\ map with SCUBA-2. The PACS 70\\,\\mic\\ image is the highest resolution image to date of the thermal emission from the system and it allows us to disentangle the emission from the warm and cold dust belts, as well as to precisely constrain the geometry of the latter ring. Specifically, we determine the inclination and PA (47\\degr\\ and 116\\degr, respectively) of the ring to within 1--2\\degr\\ and the midpoint radius of the ring ($\\approx$165\\,AU) with a precision of about 2\\%. This radius is about twice as large as would be expected for dust grains emitting like blackbodies, indicating that significant amounts of relatively small grains are present in this belt. This is in contrast with the finding that the sub-millimeter spectral index is consistent with the Rayleigh-Jeans tail, indicative of large dust grains (mm-sized and larger). We can only place an upper limit on the ring width, namely 75\\,AU at the 2$\\sigma$ confidence level, because of the still limited resolution of {\\it Herschel}. In addition, we find marginal evidence for azimuthal variations in the outer ring, as well as a possible offset of the center of the outer dust belt relative to the central star. While future higher resolution observations with ALMA are necessary to confirm these features, we note that they could be revealing the presence of unseen planets in the system. We also present deep {\\it HST} visible and near-infrared coronagraphic images of the \\ec\\ system. The outer ring is not detected in these images and we place upper limit on its scattered light surface brightness at wavelengths ranging from 0.4 to 1.7\\,\\mic.  To interpret these observations, we construct full radiative transfer models of the system's SED and 70\\,\\mic\\ image in an effort to constrain the dust properties in both belts. Assuming a simple power law grain size distribution, we find that the minimum grain size in the outer belt is about four times larger than the blow-out size, a conclusion that is also supported by the non-detection of scattered light in our {\\it HST} images. While the gap between this minimum grain size and the blow-out size could be bridged if the grain size distribution was shallower than predicted by collisional cascade models, it is evident that some mechanism must act to remove from the outer belt grains that should not be expelled by radiation pressure alone. The available data do not allow us to determine which mechanism is most plausible. On the other hand, the dust in the inner belt has to contain substantial amounts of grains at the blow-out size and, quite possibly, even below this limit. Together with the very strong emission from the inner belt, this difference in dust properties between the two dust belts suggests that the warm dust cannot be explained by mere transport of grains from the outer belt in (e.g., via a production from eccentric planetesimals), but rather support a rare and more violent recent event. However, the time baseline between the first {\\it IRAS} detection of the infrared excess in the system and the newest {\\it Herschel} observations presented here shows that the dust in the inner belt must survive for several decades or be continuously replenished."
    },
    "1402/1402.4903_arXiv.txt": {
        "abstract": "Magnetic fields govern the plasma dynamics in the outer layers of the solar atmosphere, and electric fields acting on neutral atoms that move across the magnetic field enable us to study the dynamical coupling between neutrals and ions in the plasma. In order to measure the magnetic and electric fields of chromospheric jets, the full Stokes spectra of the Paschen series of neutral hydrogen in a surge and in some active region jets that took place at the solar limb were observed on May 5, 2012, using the spectropolarimeter of the Domeless Solar Telescope at Hida observatory, Japan. First, we inverted the Stokes spectra taking into account only the effect of magnetic fields on the energy structure and polarization of the hydrogen levels. Having found no definitive evidence of the effects of electric fields in the observed Stokes profiles, we then estimated an upper bound for these fields by calculating the polarization degree under the magnetic field configuration derived in the first step, with the additional presence of a perpendicular (Lorentz type) electric field of varying strength. The inferred direction of the magnetic field on the plane of the sky (POS) approximately aligns to the active region jets and the surge, with magnetic field strengths in the range $10\\,\\mathrm{G} < B < 640 \\,\\mathrm{G}$ for the surge. Using magnetic field strengths of $70$, $200$, and $600\\,\\mathrm{G}$, we obtained upper limits for possible electric fields of $0.04$, $0.3$, and $0.8\\,\\mathrm{V\\,cm^{-1}}$, respectively. This upper bound is conservative, since in our modeling we neglected the possible contribution of collisional depolarization. Because the velocity of neutral atoms of hydrogen moving across the magnetic field derived from these upper limits of the Lorentz electric field is far below the bulk velocity of the plasma perpendicular to the magnetic field as measured by the Doppler shift, we conclude that the neutral atoms must be highly frozen to the magnetic field in the surge. ",
        "introduction": "\\label{sec.intro}  Magnetic fields govern the plasma dynamics in the outer layers of the solar atmosphere. The magnetic energy built up by convection in and below the photosphere is transferred from the photosphere into the heliosphere. During the transfer of energy into the outer layers of the solar atmosphere, i.e., chromosphere and corona, the magnetic energy is partially converted into plasma kinetic energy, causing various active phenomena in the solar atmosphere. Thus the measurement of the magnetic field is of crucial importance to identify the mechanisms responsible for the dynamical phenomena in the solar atmosphere.  Several processes can generate polarization in spectral lines in response to the presence of magnetic and electric fields in the radiation emitting plasma. The Zeeman and Stark effects are produced by the energy separation of the atomic levels into multiple sublevels due to the potential energy of the radiating atom in the external fields. As the splitted components have different polarization properties, this results in a wavelength dependent polarization across the spectral line. In the absence of velocity or field gradients in the emitting plasma, the broadband (i.e., wavelength integrated) polarization from the Zeeman and Stark effect vanishes. In contrast, in the presence of \\emph{atomic polarization} (i.e., population imbalances and quantum interference among the magnetic sublevels; e.g., \\citealt{landi04}% ), the emitted (scattered) radiation is typically characterized by a non-vanishing broadband polarization. The two dominant characteristics of atomic polarization are atomic \\emph{alignment} and \\emph{orientation}. Atomic alignment is a state of population imbalances between magnetic sublevels with different absolute values of their magnetic quantum numbers. Atomic orientation is characterized by population imbalances between sublevels with positive and negative magnetic quantum numbers. The effects of atomic alignment and orientation in the scattered radiation are observed in the form of broadband linear and circular polarization, respectively. In the process of radiation scattering, the anisotropy of the incident radiation field produces atomic alignment, and the scattered light is therefore linearly polarized, with a magnitude depending on the scattering geometry. The presence of a magnetic or electric field produces a relaxation of the atomic coherence among the magnetic sublevels resulting in a modification of the polarization (typically, depolarization and rotation) of the scattered radiation, known as the Hanle effect. Strong magnetic or electric fields also cause a conversion of atomic alignment into atomic orientation when the energy splitting of the sublevels induced atomic level crossings \\citep[A-O mechanism, ][]{lehmann64,landi82,kempetal84,favatietal87,casini05b}. Collisions of the scattering atoms with free electrons and protons tend to reduce the atomic level polarization. For the typical density of the upper solar chromosphere, collisional transitions between the fine-structure levels pertaining to the same Bohr level of the neutral hydrogen may play a significant role in the depolarization process \\citep{bommier86,sahal96,stepan11}.  Recent progress in the modeling of scattering polarization \\citep{bueno02, ariste02, bueno07} and in the development of increasingly sensitive spectropolarimetric instrumentation \\citep[e.g.][]{collados07, jaeggli10, anan12} enables us to measure the chromospheric magnetic field by interpreting the observed polarization in terms of the Zeeman and Hanle effects. Recently, the magnetic field of some types of chromospheric phenomena have been measured, for instance, in filaments \\citep{bueno02}, prominences \\citep{casini03, merenda06, sasso11}, spicules \\citep{lopez_casini05, bueno05}, and in regions of emerging magnetic flux \\citep{lagg04}.  Chromospheric jets called ``surges'' often occur near sunspots, in the proximity of neutral points of the magnetic fields \\citep{rust68}, in flux emergence regions \\citep{kurokawa88}, or in association with magnetic flux cancellation \\citep{gaizauskas96}. They reach heights up to 200 Mm and velocities of 50 to 200 km s$^{-1}$ \\citep{roy73}. Recent observations have found tiny chromospheric jets outside sunspots in active regions \\citep{shibata07, morita10}; their typical length and velocity are 1 to 4 Mm and 5 to 20 km s$^{-1}$, respectively \\citep{nishizuka11}. In addition, jet-like features are observed in sunspot penumbrae where stronger, inclined magnetic fields and weaker, nearly horizontal magnetic fields co-exist at small spatial scales \\citep{katsukawa07}. These jets are interpreted as the result of magnetic energy release caused by a change of magnetic field topology \\citep[e.g.][]{shibata07, ryutova08, nakamura12, takasao13}. However, there are no studies reporting direct measurements of the magnetic field of jets in the chromospheric lines.    Unlike the measurement of magnetic fields, the study of electric fields in solar plasmas has been given little attention. This is partially due to the fact that appreciable quasi-static electric fields on macroscopic spatial scales are unlikely in the solar atmosphere, since the relaxation time for the electric charge is very short in the highly conductive solar atmosphere. Electric fields generated in magnetic reconnection events occur at spatial scales far smaller than the spatial resolution of current spectropolarimetric instrumentation. On the other hand, the spatial scale of Lorentz electric fields associated with plasma motions across the magnetic field can be large enough to be resolved with existing solar telescopes, and the associated Stark effect can be detected in a weakly ionized plasma that is able to reach sufficient bulk velocities across the magnetic field.  Since detection of the electric field experienced by neutral atoms moving across the magnetic field can tell us about the degree of dynamical coupling between the neutrals and the ionized plasma, the observation and detection of these \\emph{motional} electric fields is of particular interest for understanding the dynamics of partially ionized plasmas in the chromosphere.  In order to possibly detect electric fields, we carried out spectropolarimetric observations of the high Paschen series (i.e., large principal quantum number of the upper level) of neutral hydrogen. The choice of hydrogen was motivated because the linear Stark effect (a first order perturbation of the atomic level energy) only occurs in hydrogen-like ions, whereas multi-electron atoms tend to show only quadratic (and higher order) energy effects in the presence of an external electric field. In the case of the linear Stark effect, the splitting of the energy levels is also roughly proportional to the square of the principal quantum number \\citep[e.g.][]{moran91, foukal95}, so the use of high transitions in a hydrogen line series is desirable in order to increase the magnitude of the polarization signal. \\cite{moran91} and \\cite{foukal95} estimated the electric field in post flare loops and prominences by studying the line width variation of the linear polarization profile between two orthogonal states of polarization. The implicit assumption in that study was that atomic polarization of such highly excited initial levels could be neglected since they should mainly be populated by isotropic recombination processes \\citep{casini96}. \\citet{foukal95} measured a surge that occurred in association with a solar flare and obtained a value of the electric field of $\\sim 35\\,\\mathrm{V\\,cm^{-1}}$.  Formalisms for the description of hydrogen line polarization in the simultaneous presence of magnetic and electric fields in both local thermodynamic equilibrium (LTE) \\citep{casini93} and non-LTE including atomic polarization \\citep{casini05b} have been developed. In LTE, an electric field of the order of 1 V cm${\\rm ^{-1}}$ produces a linear polarization of the order of 0.1\\% through the linear Stark effect. \\citet{casini05b} pointed out that if an electric field is present in the non-LTE scattering process, magnetic field strengths of order of 10 G can produce a significant level of atomic orientation through the A-O mechanism, comparable to that produced by a magnetic field of the order of 10$^3$ G in the absence of an electric field. This is because, in the presence of an electric field, magnetic-induced level crossing can happen between $\\Delta L=1$ levels that are separated in energy by a difference comparable to the Lamb shift, rather than by the much larger L-S coupling fine structure energy separation. In the scattering process with both magnetic and electric fields, an electric field of the order of 0.1 V cm${\\rm ^{-1}}$ is sufficient to modify the linear polarization of the observed Paschen lines by about 1\\%.  \\citet{gilbert02} evaluated the drain speed of neutral atoms across the magnetic field in a simple prominence model with a partially ionized plasma, in which the solar gravitational force balances with the frictional force proportional to the relative flow of the neutral and ionized components. Their analytical calculations in steady state show the downflow velocity of neutral hydrogen to be $\\sim 1.6\\,\\mathrm{km\\,s^{-1}}$ for a density of $10^{10} \\,\\mathrm{cm^{-3}}$, assuming a ionization fraction of hydrogen of $10^{-3}$ and a ionization fraction of helium of $10^{-2}$. If we assume a magnetic field strength of $100\\,\\mathrm{G}$, the electric field experienced by neutral atoms moving across the magnetic field is $\\sim 0.16\\,\\mathrm{V\\,cm^{-1}}$, and it is possible that the electric field modifies the polarization of the scattered radiation. In dynamic phenomena, electric fields can be stronger than in the steady state case. For instance, the neutral hydrogen in a surge is ejected with a velocity of 50 to 200 km s$^{-1}$ with a whiplike motion \\citep{roy73, nishizuka08}, while neutral hydrogen in prominences descends and rises along vertical threads with a velocity of 5 to 20 km s$^{-1}$ \\citep{zirker98,berger08}.  In order to study the magnetic and electric fields in chromospheric jets, we observed the full Stokes spectra of the Paschen series of \\ion{H}{1} in active region jets that took place on the solar limb on May 5, 2012. For convenience, throughout this paper we use the two terms ``surge'' and ``jet'' to refer to large and small events, respectively. In the following sections, we describe the details of the observations (Sec. \\ref{sec.obs}), the inference of the magnetic and possible electric fields (Sec. \\ref{sec.res}), and finally we discuss our measurements and provide our conclusions (Sec. \\ref{sec.dis_sum}). ",
        "conclusions": "\\label{sec.dis_sum}  We presented magnetic field inferences in a surge and in active region jets, and a first estimation of the upper limit of motional electric fields, and the corresponding limit to plasma velocities across the magnetic field lines.  The direction of the magnetic field on the POS is found to approximately align to the jets and the surge. This confirms the common scenario of the chromospheric jet model, in which plasma is ejected along the magnetic field \\citep[e.g.][]{shibata07,takasao13}. When we look carefully at the geometric relation between the magnetic field and the jet, we find that the direction of the magnetic field is slightly tilted counterclockwise from the direction of the surge. Further study with high-accuracy spectro-polarimetric data for a number of similar events is needed in order to clarify the details of the geometric relation between the magnetic field and jets.  The strength of the magnetic field in the three jets is found to be in the ranges $60\\,\\mathrm{G} < B < 630 \\,\\mathrm{G}$, $190\\,\\mathrm{G} < B < 960 \\,\\mathrm{G}$, and $40\\,\\mathrm{G} < B < 200 \\,\\mathrm{G}$. This determination is driven mainly by the observed circular polarization signal and its symmetry characteristics. Because NCP can also be generated by the coupling of velocity and magnetic field gradients along the LOS \\citep{illing74}, and since our model does not contemplate the possibility of a multi-component atmosphere, we cannot exclude the possibility that the inferred magnetic field strengths may be affected by the possible presence of velocity and magnetic field inhomogeneities along the LOS.  The upper limit of the motional electric field is estimated to be $0.04\\,\\mathrm{V\\,cm^{-1}}$ for a magnetic field strength of $70\\,\\mathrm{G}$, based on the observed polarization degree in the P10 line. If we neglect the magnetic field and the atomic level polarization, as assumed in previous estimates of solar electric fields, the limit to the electric field that is derived from the observed polarization degree (using, e.g., the formalism of \\citealt{goto08}) is $\\sim 30\\,\\mathrm{V\\,cm^{-1}}$, which is significantly larger than our estimates.  In our inversion model, we neglected collisional coupling among the atomic levels, which tends to reduce the amount of linear polarization and NCP in the scattered radiation. In the case of the surge, the observed linear polarization may be explained through a combination of collisional depolarization with the polarizing effects of the magnetic field and the possible motional electric fields in the plasma (Fig.~\\ref{fig.e_pd}). In particular this implies that our inferred upper bounds for the motional electric field present in the surge could be underestimated. Similarly, our model might underestimate the LOS component of the magnetic field, if the observed low values of NCP were caused by the destruction of atomic orientation by collisions. Finally, we note that the presence of isotropic collisions does not cause a significant change of the direction of the emergent linear polarization, hence we can still reliably determine the azimuth angle of the magnetic field on the POS ($\\Phi_{B}$) from the inversion, independently of the effects of collisions.    From the upper limit of the motional electric field, we can estimate the maximum velocity of neutral hydrogen moving across the ambient magnetic field. At 3:16 UT, the surge reaches its maximum height (Fig.~\\ref{fig.v}) and a large bulk velocity (i.e., Doppler shift) is observed (Fig. \\ref{fig.scatres}). Therefore the velocity vector can be assumed to be along the LOS and perpendicular to the magnetic field (Fig.~\\ref{fig.scatres}). On the other hand, from our analysis of the polarization effects of motional electric fields, the upper limit of the velocity of neutral hydrogen moving across the magnetic field is found to be $1.1\\,\\mathrm{km\\,s^{-1}}$ in this geometry. As this is much smaller than the measured bulk velocity of $12.5\\pm1.4\\,\\mathrm{km\\,s^{-1}}$, our estimate of the upper bound of the motional electric fields in the plasma leads us to conclude that neutrals must be in a highly frozen-in condition in the surge."
    },
    "1402/1402.6906_arXiv.txt": {
        "abstract": "The Gemini Multi-conjugate Adaptive Optics System - GeMS, a facility instrument mounted on the Gemini South telescope, delivers a uniform, near diffraction limited images at near infrared wavelengths (0.95 $\\mu$m - 2.5 $\\mu$m) over a field of view of 120$\\:$\\arcsec. GeMS is the first sodium layer based multi laser guide star adaptive optics system used in astronomy. It uses five laser guide stars distributed on a 60$\\:$\\arcsec square constellation to measure for atmospheric distortions and two deformable mirrors to compensate for it. In this paper, the second devoted to describe the GeMS project, we present the commissioning, overall performance and operational scheme of GeMS. Performance of each sub-system is derived from the commissioning results. The typical image quality, expressed in full with half maximum, Strehl ratios and variations over the field delivered by the system are then described. A discussion of the main contributor to performance limitation is carried-out. Finally, overheads and future system upgrades are described. ",
        "introduction": "\\label{sec:introduction}  Adaptive Optics (AO) is a technique that aims to compensate the phase aberrations induced by atmospheric turbulence. Aberrations are measured by a Wave-Front Sensor (WFS), using observations of a Guide Star (GS). Corrections are applied by an optical active device, generally a Deformable Mirror (DM). For the current class of 8-10 meter astronomical telescopes, AO typically improves the angular resolution by an order of magnitude, and restores a resolution close to the telescope diffraction limit. Over the past 20 years, AO for astronomy has gone from a demonstration phase to a well-proven and operational technique, and it is now almost universally considered as an essential part of any new large telescope. In addition, to increase the number of targets on which AO can be used, all of the major 8 meter telescopes are now equipped with Laser Guide Stars (LGS - see e.g. \\cite{wizinowich2012progress}). AO and LGS-AO observations have enabled major discoveries in astronomy with, among others, the discovery and study of the supermassive black hole at the center of our Galaxy (e.g. \\cite{ghez2008gc, genzel2010gc}), detailed images of the surface of solar systems bodies (e.g. \\cite{hartung2004new, depater2010persistent}), or precise morphology and dynamics of very distant galaxies (e.g. \\cite{huertas2008robust, cresci2009sins, wright2009dynamics, carrasco2010}).  The advent of a new generation of AO systems called Wide Field AO (WFAO) mark the beginning of a new era. By using multiple GSs, either LGS or Natural Guide Stars (NGSs), WFAO significantly increases the field of view of the AO-corrected images, and the fraction of the sky that can benefit from such correction. Where the first AO systems (also called Single-Conjugate AO or SCAO) were well suited for observations of bright and relatively compact objects, the new generation of WFAO is opening the path for a multitude of new science studies.  Different flavours of WFAO have been studied over the past years. They all require multiple GSs to perform a tomographic analysis of the atmospheric turbulence. What differentiates the various WFAO systems is how the turbulence correction is applied. Ground Layer AO (GLAO) uses a single DM optically conjugated to the telescope pupil \\citep{rigaut2001glao}. If the correction is optimised over a field of view larger than the anisoplanatism angle, then only the atmospheric layers close to the ground will be compensated \\citep{ragazzoni2002multiple}, providing a partial, but uniform correction over the field.  Another solution, called Multi-Conjugate AO \\citep[MCAO,][]{dicke1975phase,beckers1988increasing,ellerbroek1994first,johnston1994analysis} uses several DMs optically conjugated to the main turbulence layers. In that case, all the layers close to the DM altitude conjugation will be compensated, restoring the telescope diffraction limit over field of views many times larger than the ones achievable with SCAO at Near Infra-Red (NIR) wavelengths.   MCAO for night time astronomy\\footnote{MCAO systems for solar astronomy have been in use since the mid 2000s at the VTT in Tenerife and at the Dunn solar telescope at Sacramento Peak.} was first demonstrated by MAD, a prototype built at the European Southern Observatory \\citep{marchetti2003mad,marchetti2007mad}. MAD used three NGSs, two DMs conjugated at the ground and at an altitude of 8.5$\\:$km, and provided a corrected field of view of almost 2$\\:$arcmin across. Although MAD successfully demonstrated the gain brought by WFAO over SCAO, it was limited in the number of potential targets due to limiting magnitude of the required NGSs (m$_{\\rm R}<12.5$) and, essentially running out of targets, the instrument was decommissioned in 2008. The first multi-LGS WFAO system open for the community was a GLAO system operating at the MMT, which uses three 532$\\:$nm Rayleigh LGSs \\citep{Baranec2009glao}. GeMS, the Gemini MCAO system, is the first sodium based multi-LGS MCAO system.  This paper is the second of a review describing the GeMS project. The first paper (\\cite{rigaut2013review} - hereafter Paper I) covers the first part of the history of the project, from the original idea to the first light images. It also includes a detailed description of GeMS, hence only a brief description of the system is given here. GeMS is made by two main sub-systems: (i) the LGS Facility (LGSF) that includes a 50$\\:$W laser and an optical system called Beam Transfer Optics (BTO) that relays the laser light, and controls the LGSs, and (ii) the MCAO bench, called {\\sc Canopus}. In short, the 50$\\:$W laser is split in 5$\\times$10$\\:$W beams to produce the 5 LGSs projected on the sky at the corners and center of a 60\\arcsec square. These LGSs feed five 16$\\times$16 subapertures Shack-Hartmann WFSs (so-called LGSWFSs). The 2040 slope measurements are used to compute the MCAO high-order correction, correction provided at up to 800$\\:$Hz by two deformable mirrors conjugated to 0 and 9$\\:$km. In addition, up to three visible NGSs provide the measurements for the compensation of the tip-tilt and anisoplanatic modes. The tip-tilt compensation is done by a tip-tilt mirror (TTM) while the Tilt-Anisoplanatic (TA) modes are compensated by a combination of quadratic modes on DM0 and DM9. A fraction of the light from one of the NGS is directed toward a Slow Focus WFS (SFS), which controls the LGSWFS zoom to keep the instrument in focus. At the GeMS output, the corrected beam can be steered toward different science instruments attached to the Cassegrain focus instrument cluster. The main instrument used to date is GSAOI \\citep{mcgregor2004gemini}, a 4k$\\times$4k NIR imager covering $85''\\!\\times 85''$ designed to work at the diffraction limit of the 8-meter telescope. %  This paper focuses on the commissioning, overall performance and operation scheme of GeMS. The goal of this paper is to give a top-level view of the GeMS capability, that could be used for instance when preparing observations. Section~\\ref{sec:commissioning} summarises the commissioning period, and details the performance of the sub-systems. Section~\\ref{sec:system_verification} gives an overview of the System Verification (SV) period, and illustrates the science capability provided by GeMS. Section~\\ref{sec:performance} analyses the top-level performance delivered by GeMS in term of image quality over the field and astrometry precision. Note that this paper does not intend to perform a detailed analysis of the system performance, as this will be presented in a dedicated paper. Section~\\ref{sec:operations} discusses the operational scheme of GeMS, including overheads, and finally section~\\ref{sec:the_future_of_gems} presents the system upgrades. ",
        "conclusions": "\\label{sec:conclusion2} GeMS has proven to be an exciting and challenging project. Because many complex subsystems must work in unison, GeMS imposes a level of complexity rarely found in an observatory facility instrument. The amount of time on sky (about 100 nights) for commissioning activity and work towards the transition into regular operations is testament alone to the complexity of the system. Operating such a complex instrument is also a challenge for the Gemini observatory and a dedicated effort had been put in place in order to simplify and stream-line operations. %  GeMS is also a pathfinder instrument, paving the way for the future developments of such systems. It is important to emphasize that no show stoppers have been found on the way, although the delivered performance differs from what was predicted during the design phases. The compensation performance is currently limited by servo lag, noise and generalised fitting, but other, and unexpected terms are also affecting the performance. Furthermore, several calibrations designed during the early phase of the project proved to be either not working or sub-optimal during the commissioning, and new solutions had to be developed. As discussed in this paper, work is continuing on the system to improve on all these aspects, as well as many others.  But if GeMS is complex, is it also a rewarding instrument. The programs executed to date and associated images have already demonstrated its huge and unique scientific potential. Images with FWHM of 0.08$\\:$\\arcsec or better over the 85$\\:$\\arcsec$\\!\\times$ 85$\\:$\\arcsec field of view are typically obtained under median seeing or better."
    },
    "1402/1402.2378_arXiv.txt": {
        "abstract": "M type stars are good targets in the search for habitable extrasolar planets. Because of their low effective temperatures, the habitable zone of M stars is very close to the star itself. For planets close to their stars, tidal heating plays an important role in thermal and orbital evolutions, especially when the planet orbit has a relatively large eccentricity. Although tidal heating interacts with the thermal state and orbit of the planet, such coupled calculations for extrasolar planets around M star have not been conducted. We perform coupled calculations using simple structural and orbital models, and analyze the thermal state and habitability of a terrestrial planet. Considering this planet to be Martian sized, the tide heats up and partially melts the mantle, maintaining an equilibrium state if the mass of the star is less than 0.2 times the mass of the Sun and the initial eccentricity of the orbit is more than 0.2. The reduction of heat dissipation due to the melted mantle allows the planet to stay in the habitable zone for more than 10 Gyr even though the orbital distance is small. The surface heat flux at the equilibrium state is between that of Mars and Io. The thermal state of the planet mainly depends on the initial value of the eccentricity and the mass of the star. ",
        "introduction": "The search for extrasolar planets has revealed that there are many terrestrial planets out of our solar system (e.g., L\\'{e}ger et al. 2009). By detecting these extrasolar terrestrial planets, many researches have tried to determine the planets' thermal and orbital histories (Henning et al. 2009; Barnes et al. 2010; L\\'{e}ger et al. 2011; Wagner et al. 2012). The thermal and orbital states of extrasolar terrestrial planets are important for the planets' evolution. In addition, the thermal and orbital states affect the habitability, making analysis of these factors of key importance.  Recently, in the search for habitable extrasolar planets, M stars were suggested for observation (Scalo et al. 2007) because the number of M stars is much larger than that of G stars. Conventionally, M stars are not considered to be likely hosts of habitable planets because their planets' orbits are too close, which results in the planets being tidally rocked (synchronously rotated) and their atmosphere freezing on the dark side. However, Joshi et al. (1997) hypothesized that carbon dioxide can make a planet habitable. Thus, the importance of M star extrasolar planet analysis continues to increase.  Focusing on the terrestrial planets in our solar system, the main heat sources are the insolation from the Sun and the radioactive decay in the planets themselves. However, if extrasolar planets orbit around small stars such as an M star, the tidal heating has a significant effect on the habitability. Compared with the Sun, M stars have low temperatures, and thus the magnitude of the insolation becomes low. Therefore, the habitable zone is closer to the star. The rate of change of a planet's orbit due to tidal dissipation increases with decreasing orbital distance. For this reason, a planet orbiting an M star may not be able to stay within the habitable zone for a long time. In addition, assuming that the planets orbit according to Kepler's law, the tidal frequency is high, producing tidal heating that may exceed the radiogenic and insolation heating. In contrast to the terrestrial planets in our solar system, such tidal heating is an important factor for the thermal and orbital evolutions of planets orbiting M stars.  The effects of tidal heating on terrestrial planets orbiting M stars have been analyzed (Jackson et al. 2008; Barnes et al. 2008). Although these researches revealed that tidal heating influences a planet habitability, they assumed that the magnitude of the dissipation of the planet is constant. The tidal heating mechanism is complicated and interacts with the internal thermal state and orbit of the planet (e.g., Hussmann \\& Spohn 2004). When a planet is tidally heated, structure and thermal states of the planet changes, which affects the tidal heating by the feed-back. In addition, the orbit and heating rates also interact with each other. Tidal heat dissipation is dependent on the eccentricity of the orbit; the eccentricity decreases with  heating rate because orbital energy is lost as tidal heat. Thus, tidal heating should be coupled with orbital and thermal evolution.  In this work, we evaluate the coupled calculation of tidal heating considering the thermal and orbital states. However, coupled calculations are very complex. In addition, the internal structure of extrasolar terrestrial planets is not constrained well.  Thus, as a fist step, we consider a Martian-sized terrestrial planet. An advantage of this choice is that the internal pressure of planets of this size is small compared with earth-sized planets, due to their small radii, which allows a simple model of the internal structure to be used. A Martian-sized planet has been discovered (Muirhead et al. 2012); however, this planet is outside the habitable zone. Despite this, we can expect similarly sized planets to be discovered within habitable zones in future observations. Considering the observation, Earth-sized planets are easier to detect compared to Martian-sized planet. However, we try the coupling theoretical calculation with Martian structure. In our solar system, Io is a tidally locked rocky satellite with high tidal heating (Segatz et al. 1988). Thus, we reference the thermal model of Io as well as Mars. If accurate internal structures of super-earths are derived in the future, we can apply the results of this work to such models with more complicated structures. ",
        "conclusions": "\\subsection{Morphology} If the tidal heating is sufficiently large, the mantle is heated and the surface heat flux exceeds the Martian heat flux. In this work, we assume stagnant lid convection as a heat transport mechanism. As an alternative  mechanism, plate tectonics may occur in planets in a hot state. Although it is currently only possible to observe the plate tectonics of Earth, it has been suggested that plate tectonics occurs in large, rocky extrasolar planets (e.g., Valencia et al. 2007). However, plate tectonics is a very complicated system and the main induction mechanism is not well understood. The size of the planet may affect the stress (Valencia et al. 2007; O'Neill \\& Lenardic 2007); however, Korenaga (2010) suggests that the effect of surface water is more important than the size of the planet. In our model, we do not include surface oceans and the size of the planets considered is relatively small.  Thus, stagnant-lid convection is more reasonable. However, in the case of Mars, plate tectonics has been suggested to be the origin of the early internal dynamo (Nimmo \\& Stevenson 2000). Thus, plate tectonics cannot be ruled out for Martian-sized extrasolar planets.  Delamination may be more probable in small-sized extrasolar planets than plate tectonics. In a hot state, the melt in the mantle reduces the mantle viscosity. Thus, the bottom of the stagnant lid is delaminated and absorbed into the mantle. Delamination displaces the surface of the planet, which is the suggested cause of the corona of Venus (Smrekar \\& Stofan 1997).  Thus, we can suggest that the planet is heated by tides, and the morphology of the planet is mainly a product of localized delamination processes caused by partial melting. As an alternative heat transfer model, heat pipe mode is suggested in Io and the Earth (Moore 2001; Moore \\& Webb 2013), which can be the important heat transfer mechanism for the tidally heated extrasolar planets. Thus, the effects of delamination and heat pipe should be considered in future work.    \\subsection{Magnetic field} It was revealed that Mars had an intrinsic magnetic field in the past (Acu\\~{n}a et al. 1999). The most probable origin of the intrinsic magnetic field is the dynamo in the core. The mechanism of the Martian dynamo is not well understood. However, it is required that more heat than that transported by adiabatic heat flux is transported from the core to the mantle in order to induce the core convection (Nimmo \\& Stevenson 2000).  If tidal heat is not generated in the mantle, heat is transported sufficiently by the mantle and a dynamo may be induced. However, in our calculations, tidal heating heats up the mantle. In some cases, heat is even transported from the mantle to the core. Thus, in the hot state of a planet, it is unlikely the dynamo is activated.  An intrinsic magnetic field is not predicted in a tidally heated planet; however, if the orbital plane of the planet is oblique to the magnetic axis of the star, an induced magnetic field may be possible because of the melt in the mantle. In the case of Io, an induced magnetic field is observed (Khurana et al. 2011) . The induced magnetic field of Io is consistent with a mantle of at least 20 \\% partial melt (Khurana et al. 2011). This melt fraction is comparable to our calculation results.  \\subsection{Heterogeneity of the planet} For simplicity, we assume that the internal structure of the planet is spherical shells and depends only on the radial distance. As for the habitable zone, the spherical-shell structure should provide a valid approximation because the temperature in light side is not so high. However, the heterogeneity of the planet is an interesting topic for future study, especially for planets that orbit significantly close to their stars. Tidal heating in a planet with a heterogeneous structure was calculated for the case of Enceladus, though it is extremely complicated (Tobie et al. 2008; B\\v{e}hounkov\\'a et al. 2012). If coupled calculations including a heterogeneous structure are possible in future work, an important comparison to the spherical-shell model of this work can be made.  \\subsection{Resonance} One caveat is that the coupled calculations in this work do not consider the resonance of the planet. If the planet is in a resonant state with other planets, the orbital evolution becomes significantly more complicated. Resonance among planets increases the planets' eccentricity, as exhibited by the Jovian satellites Io, Europa and Ganymede (Murray \\& Dermott 1999). Some satellites around Saturn also interact with each other (Murray \\& Dermott 1999). Hence, if there are multiple planets around a star, we must consider the resonance between them."
    },
    "1402/1402.5066_arXiv.txt": {
        "abstract": "We present maps of a large number of dense molecular gas tracers across the Central Molecular Zone of our Galaxy. The data were taken with the CSIRO/CASS Mopra telescope in Large Projects in the 1.3\\,cm, 7\\,mm, and 3\\,mm wavelength regime. Here, we focus on the brightness of the shock tracers SiO and HNCO, molecules that are liberated from dust grains under strong (SiO) and weak (HNCO) shocks. The shocks may have occurred when the gas enters the bar regions and the shock differences could be due to differences in the moving cloud mass. Based on tracers of ionizing photons, it is unlikely that the morphological differences are due to selective photo-dissociation of the molecules. We also observe direct heating of molecular gas in strongly shocked zones, with a high SiO/HNCO ratios, where temperatures are determined from the transitions of ammonia. Strong shocks appear to be the most efficient heating source of molecular gas, apart from high energy emission emitted by the central supermassive black hole Sgr A* and the processes within the extreme star formation region Sgr B2. ",
        "introduction": "The Central Molecular Zone (CMZ) covers the inner $\\sim 500$\\,pc of the Milky Way and contains a substantial amount of molecular gas (several $10^{7}$\\,M$_{\\odot}$ e.g. \\cite[Oka et al. 1998]{oka98}). Models show that the gas may be funneled along the central Galactic bar into the CMZ and eventually accretes on $x_2$ orbits that cover the inner $\\sim 200$\\,pc (e.g. \\cite[Morris \\& Serabyn 1996]{mor96}). Eventually the gas will be destroyed by environmental influences, during the process of star formation, or, in rare cases, the dense gas will be transported all the way further toward Sgr A*, the central supermassive black hole of the Milky Way.  In this paper, we focus on material that is formed during shocks. HNCO, for example, is a molecule that is formed on dust grains and liberated into the gas phase by weak shocks (\\cite[Meier \\& Turner 2005]{mei05}). Stronger shocks can be traced by SiO where the silicon is provided by shock destruction of the dust grains themselves (\\cite[Martin-Pintado et al. 1997]{mar97}). Since the ISM conditions in the Galactic Center are dominated by shock physics, we use these two molecules to trace the environment as the gas accumulates in the CMZ. \\cite{jon12} presents maps of both molecules using the ATNF/CASS Mopra\\footnote{The Mopra radio telescope is part of the Australia Telescope National Facility which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. } telescope and \\cite{ott08} shows single dish Mopra maps of ammonia, a molecule which is used as a thermometer to trace shock heating. ",
        "conclusions": ""
    },
    "1402/1402.0933.txt": {
        "abstract": "Magnetic fields of   laboratory, planetary, stellar, and galactic plasmas commonly  exhibit significant order on  large temporal or spatial  scales  compared to the otherwise random motions within the hosting system.  Such ordered fields can be  measured in  the case of planets, stars, and galaxies, or inferred indirectly by the action of their dynamical influence, such as  jets.  Whether large scale fields are amplified in situ or a remnant from previous stages of an object's history is often  debated for objects without  a definitive magnetic activity cycle.  Magnetic helicity, a measure of twist and linkage of magnetic field lines, is a unifying  tool for understanding  large scale field evolution for both mechanisms of origin. Its importance  stems from its two basic properties:  (1)  magnetic helicity is typically  better conserved  than magnetic energy; and  (2) the magnetic energy associated with a fixed amount of magnetic  helicity is minimized when the system relaxes  this helical structure to the largest scale available. Here I  discuss how magnetic helicity has come to  help us understand the saturation of and sustenance of  large scale dynamos, the need for either local or global helicity fluxes to avoid dynamo quenching, and the associated observational consequences. I  also discuss how  magnetic helicity acts as a hindrance to turbulent diffusion of large scale fields,  and thus a helper for fossil remnant large scale field origin models in some contexts.  I  briefly discuss the connection between large scale fields and accretion disk theory as well. The goal here  is to provide a conceptual primer to help the reader efficiently  penetrate the  literature. ",
        "introduction": "Planets,  stars,   galaxies  are all examples of astrophysical rotators that reveal direct or indirect evidence for large scale ordered magnetic fields (Schrijver  \\&  Zwaan 2000; Brandenburg \\& Subramanian 2005; Shukurov 2005; Beck 2012 Roberts \\& King 2013). Here ``large\" scale implies a coherent flux on scales comparable to the size of the hosting rotator and most importantly, larger  than the scale of fluctuations associated with  chaotic turbulent flows.   In fact,  all of these systems show evidence for both small and large scale magnetic fields, so the fact that large scale order persists even amidst a high degree of smaller scale  disorder is  a  core challenge of explaining the emergence of large scale magnetic structures across such disparate classes of  rotators.  Large scale magnetic fields are also likely fundamental to coronae and jets from accretion engines  around young and dying stars and compact objects (e.g. Blandford \\& Payne 1982; K\\\"onigl 1989; Field \\& Rogers 1993; Blackman et al. 2001;  Lynden-Bell 2006; Blackman \\& Pessah 2009; Pudriz et al. 2012; Penna et al. 2013). Accreting systems typically exhibit a continuum spectrum and luminosity best explained by matter accreting onto a central object falling deeper into a potential well and thereby releasing positive kinetic energy in the form  of radiation or jet outflows. In fact, as  Fig. {\\ref{bernfig1} shows,  the  classes of objects  likely harboring jets has increased  as observations have improved, and jets likely indicate the role of large scale magnetic fields.  In young stars,  pre-planetary nebulae, micro quasars, and active galactic nuclei the jets typically  have too much collimated momenta to be driven by  mechanisms that do not involve  large scale magnetic fields (Bujarrabal et al. 2001; Pudritz et al. 2012).  From the jets of AGN, Faraday rotation from ordered helical magnetic fields is directly observed (Asada et al. 2008; Gabuzda et al. 2008; Gabuzda et al.  2012).  Because jets are anchored in the accretion engines, they play a role in extracting the angular momentum that allows remaining disc material to accrete.  The ionization fractions of accretion disks are commonly high enough, at least in regions near the very center,  to be unstable to the magneto-rotational instability (MRI) (Balbus \\& Hawley 1991, 1998, 2003). A plethora of numerical simulations now commonly reveal  that   the systems evolve to a nonlinear turbulent steady state whose Maxwell stress dominates the Reynolds stress and  for which  large scale ordered magnetic  fields emerge with  cycle periods  $\\sim 10$ orbit periods (e.g. Brandenburg et al. 1995; Davis et al. 2010; Simon et al. 2011; Guan \\& Gammie 2011; Sorathia et al. 2012; Suzuki and Inutsuka 2013). With the caveat that  angular momentum plays a comparatively  subdominant role in structural support for stars, the coronae of stars and the emergence of large scale ordered fields that thread coronal holes  of the sun (e.g. Schrijver  \\&  Zwaan 2000) can help shed light on  related phenomena of the rising and opening up of large scale fields that  form jets and coronae from accretion disks. % (REF blackman 2). And Both types of coronae likely have a full spectrum of magnetic field scales.  There are three possibilities for the  origin of large scale fields in astrophysical rotators:  The first is that the contemporary field is simply the result of advection and compression of the field that was present in the object before it was formed (e.g. Braithwaite \\& Spruit 2004; Kulsrud et al. 1997;  Lovelace et a. 2009; Subramanian 2010;  Widrow et al. 2012). The second is that the field is in fact dynamo produced in situ, extracting free energy from shear, rotation and turbulence in such  a way as to  sustain the field against turbulent diffusion (e.g. Moffatt 1978;  Parker 1979; Ruzmaikin et al. 1988; Shukurov  2005; Charbonneau 2013).  The third possibility is  some combination of the two (e.g. Kulsrud\\& Zweibel 2008). In systems like the Sun and Earth where cycle periods involving field reversals are observed, the need for in situ dynamo amplification is unambiguous. For galaxies or accretion disks, the  evidence for  in situ amplification of large scale fields is more indirect. Regardless of  whether the fields are initially frozen in and advected, amplified in situ, or a combination of the two, the evolution of   magnetic helicity  is very helpful for  understanding the physics of magnetic field origin as we shall see.  Much of what we can observationally infer about  magnetic fields of planets, stars, and accretion disks comes from information external to where the real action of magnetic field amplification and conversion of kinetic to magnetic energy occurs. Yet, most theory and simulation of astrophysical dynamos focuses on the interiors.   This  situation contrasts that of our own galaxy where we observe the field from within (e.g. van Eck et al 2011; Beck 2012), albeit on time scales too short to observe its dynamical evolution.   A related  point  is that  the hidden interiors of astrophysical rotators are typically flow-dominated, with the magnetic field energy density generally weaker than that of the  kinetic energy.  However in the surrounding coronae of stars, accretion disks (and maybe even for galaxies) the field  dominates the kinetic energy. Thus we must learn about the flow-dominated interiors from observations of the magnetically dominated exteriors and understand the coupling between the two. Laboratory plasmas of fusion devices are in fact magnetically dominated and there are many lessons learned from this contexts.   Tracking magnetic helicity evolution has proven helpful for understanding the field evolution in both magnetically dominated and flow-dominated circumstances.  In this paper  I discuss basic principles of magnetic helicity evolution  to  guide the physical intuition of how large scale magnetic fields arise and evolve.  This overview represents one path through the subject and is  intended as a conceptual primer to ease immersion in the literature rather than a  complete detailed  review.  Many relevant  papers  will therefore regrettably go uncited---a situation that I  find increasingly difficult to avoid. A example of an earlier  more detailed review  is Brandenburg \\& Subramanian (2005).  In section 2,  I discuss the key physical   properties of magnetic helicity that are central to all subsequent topics discussed. In section 3, I summarize the conceptual progress of how these principles apply to modern developments in magnetic dynamo theory. In section 4, I describe the simple  two-scale model  for closed systems that has come to be  useful for gaining unified insight  of large scale field evolution and dynamo saturation more quantitatively. I discuss several  applications of the two-scale theory: dynamo saturation,  the  resilience of helical fields to turbulent diffusion, and  dynamical relaxation. In section 5, I discuss the role of helicity fluxes and their relation to the results  section 4. In section 6, I briefly discuss the issue of gauge non-invariance of magnetic helicity. I conclude in section 7, emphasizing the expectation that  both signs of magnetic helicity on different scales should appear in astrophysical rotators with coronal cycle periods, and comment on the connections between magnetic helicity dynamics and accretion disk theory. ",
        "conclusions": "Tracking magnetic helicity helicity in MHD systems  is  an important, unifying tool to understand the the processes by which large scale magnetic fields form and evolve in both astrophysical and laboratory plasmas. The purpose of this review has been to provide one  path through the forrest as a conceptual primer to the literature.  Three key principles underlie the role of magnetic helicity in all contexts: (1) magnetic helicity is is a measure of twist, linkage or writhe; (2) magnetic helicity is better conserved than magnetic energy under most circumstances for a closed system with or without velocity flows; (3) the energy in a helical magnetic field is minimized when the field relaxes to the largest scale available, consistent with the boundary conditions.  The importance of magnetic helicity for large scale field generation is evident from early studies of helical MHD turbulence (Pouquet et al. 1976), but incorporating its role into a 21st century dynamical mean field theory has emerged only in the pas decade or so. The 20th century textbook dynamos, unlike 21st century theory, do not conserve magnetic helicity and as such were unable to predict or reveal how LSDs saturate.  A key aspect of 21st century theory is that the growth of large scale fields, facilitated by an EMF, involves growth of large scale magnetic helicity and small scale magnetic helicity of the opposite sign. In the absence of helicity fluxes,  the small scale build up suppresses further growth of large scale field. In the case of sheared rotators,  unless helicity fluxes can remove the offending small scale  magnetic helicity, the large scale field  not only saturates  but  may even decay. Alternatively expressed, it seems that astrophysical dynamos, like laboratory plasma dynamos may involve an  EMF that is commonly aided or sustained by the divergence of small scale helicity fluxes.   In a quasi-steady state for the sun, a crude minimalist prediction is that both signs of helical magnetic fields  should appear in the northern hemisphere with small scale left handed structures and right handed large scale structures with the reversed combination in the southern hemisphere. More efforts to measure the spectral distribution of helical fields in the solar corona and wind would be valuable. Large scale dynamo models for the galaxy and for  accretion disks that incorporate magnetic helicity dynamics have also been emerging.  Accretion disks pose an interestingly rich opportunity of study for helicity dynamics and large scale dynamos because traditionally large scale dynamo theory has been studied independent of theories of angular momentum transport. The ubiquity of large scale dynamos seen in simulations and the ubiquity of observed astrophysical coronae and jets indicates that a significant contribution to angular momentum transport comes from large scale fields. This needs to be incorporated into a combined mean field accretion disk theory that captures local and large scale angular momentum transport, and large scale field growth.  Coronae of stars, disks, and laboratory plasmas are all magnetically energy dominated. In magnetically dominated environments   the principles of magnetic helicity evolution have long been helpful to understand the evolution of magnetic structures subject to their foot-point motions. The helicity injection by  foot-point motions is analogous to  injection of small scale helicity in laboratory devices, where the  system responds by relaxing the helicity to large scales  In astrophysical coronae, it is likely that some contribution from both signs rather than a single sign are injected, so the relaxation process must take this into consideration globally,even if local structures are injected with primarily one sign.  It was in fact in the context of  laboratory plasma magnetic relaxation where the importance of helicity fluxes was first identified.  Finally, as reviewed herein, the basic properties of magnetic helicity also underlie its role in making large scale helical fields resilient to turbulent diffusion.  Recent work on this topic for closed systems may strengthen the potential efficacy of fossil field origin of large scale fields in some astrophysical contexts. A helical field (in the absence of global helicity fluxes) is much more resilient to turbulent diffusion than  non-helical large scale fields. The effect is best understood not as reduction of the turbulent diffusion coefficient, but rather as a competition between the unfettered turbulent diffusion and additional competing  tendency for helicity to relax back toward the largest scales. The  driver for this inverse transfer is the very small scale magnetic helicity that is sourced by initial diffusion from the large scale helicity in the first place."
    },
    "1402/1402.5120_arXiv.txt": {
        "abstract": "{We investigate the effect of non-canonical kinetic terms on inflation in supergravity. We find that the biggest impact of such higher-derivative kinetic terms is due to the corrections to the potential that they induce via their effect on the auxiliary fields, which now have a cubic equation of motion. This is in contrast to the usual (non-supersymmetric) effective field theory expansion which assumes that mass-suppressed higher-derivative terms do not affect the lower-derivative terms already present. We demonstrate with several examples that such changes in the potential can significantly modify the inflationary dynamics. Our results have immediate implications for effective descriptions of inflation derived from string theory, where higher-derivative kinetic terms are generally present. In addition we elucidate the structure of the theory in the parameter range where there are three real solutions to the auxiliary field's equation of motion, studying the resulting three branches of the theory, and finding that one of them suffers from a singularity in the speed of propagation of fluctuations.}      \\begin{document} ",
        "introduction": "The observational data available on the cosmic microwave background (CMB) \\cite{Smoot:1992td, Komatsu:2010fb, Dunkley:2010ge, Story:2012wx}, including the latest PLANCK data \\cite{Ade:2013zuv, Ade:2013uln}, baryon acoustic oscillations \\cite{Percival:2009xn}, redshift \\cite{Riess:1998cb, Perlmutter:1998np} and Hubble parameter measurements \\cite{Riess:2011yx}, portrays a universe very close to being spatially flat, in which large-scale structure was generated from an almost scale-invariant power spectrum of primordial density perturbations with Gaussian distribution. A major contemporary challenge in cosmology is to find a convincing, self-consistent, explanation for this early state of the universe. The two most promising theories to date are a period of single-field plateau-type inflation \\cite{Ade:2013uln}, and certain classes of two-field cyclic models of the universe \\cite{Lehners:2013cka,Fertig:2013kwa}. Both theories fit the data well, subject to important assumptions: in the inflationary case, one assumes for instance that the right initial conditions for inflation were present for its onset, and that, after regulating the infinities that eternally inflating plateau models lead to, the naive predictions still hold (for contrasting views on these issues, see \\cite{Ijjas:2013vea} and \\cite{Guth:2013sya}). In cyclic models, the most important assumption is that a smooth, non-disruptive bounce from the contracting to the expanding phase is possible. In the present paper, we will only look at the inflationary case, and we will assume that the assumptions mentioned above are justified. Our concern will rather be with the (classical) dynamics during the inflationary phase.  In the absence of a UV-complete theory of quantum gravity, the dynamics of inflation is usually treated using an effective field theory (EFT) approach. Any given theory of inflation is then only valid up to some energy cut-off  $\\Lambda$, with the physics above this scale being integrated out. The Lagrangian is organized into a series of kinetic and potential terms with increasing numbers of fields and/or derivatives, in such a way that successive terms are suppressed by increasing powers of the cut-off $\\Lambda/M_{P}$ (expressed in Planck units). For non-derivative interaction terms, it is well known that, despite this suppression, the predictions of inflation are highly sensitive to certain higher-order terms in the series. For instance, the underlying UV completion can lead to large corrections to the inflaton mass via dimension 6 operators of the form $\\frac{{\\cal O}_6}{M_P^2} \\sim \\frac{{\\cal O}_4}{M_P^2} \\phi^2$, when ${\\cal O}_4$ has a VEV of the order of the inflationary energy density. This is known as the $\\eta$ problem and is a generic problem in supergravity theories of inflation, see e.g. \\cite{Baumann:2009ni}. The underlying UV-complete theory can also lead to important corrections to the kinetic terms, resulting in higher-order corrections such as the square of the ordinary kinetic term, $(\\partial \\phi)^4$ \\cite{Silverstein:2003hf, Franche:2009gk}. These higher-derivative kinetic terms can have dramatic effects and have been studied extensively, for instance in DBI inflation \\cite{Silverstein:2003hf, Alishahiha:2004eh, Shandera:2006ax, Chen:2004gc, Chen:2005ad, Kecskemeti:2006cg}, in the context of the Horndeski action \\cite{Horndeski:1974wa, RenauxPetel:2011sb, Ribeiro:2012ar} or within the EFT of inflationary perturbations \\cite{Cheung:2007st, Senatore:2009gt, Creminelli:2010qf, Baumann:2011su, Achucarro:2012sm}.  Whether the correct UV-complete underlying theory of inflation is string theory or some other high-energy theory is as yet unknown. Nevertheless, one may expect this theory to be supersymmetric for various reasons including unification of the electroweak and strong forces at high energies, the hierarchy problem \\cite{Dimopoulos:1981zb, Dimopoulos:1981yj} and phenomenology \\cite{Lukas:1998yy, Braun:2005nv}. The proper low-energy setting for a theory of inflation is then supergravity, and the goal of the present paper is to investigate a number of novel properties that the extension of the inflationary EFT framework to supergravity entails. An initial study of this extension was undertaken by Baumann and Green in \\cite{Baumann:2011nm, Baumann:2011ws}, where they studied inflationary {\\it perturbations} and focussed on possible non-Gaussian observational signatures. In that context, Baumann and Green argued that they could neglect the auxiliary fields. In the present paper, we will focus on the inflationary {\\it background} -- in that situation, the auxiliary fields play a crucial role. In fact, it is the interplay of the higher-derivative kinetic terms and the auxiliary fields that leads to all of the new effects that we found, as we will briefly describe now.  The most important field in our study will be the auxiliary field $F$ of a chiral superfield $\\Phi.$ Assuming that the auxiliary fields of the supergravity multiplet have already been eliminated via their equations of motion, the formula for the potential reduces to the expression \\be V=-e^{K/3} K_{,AA^\\star}FF^\\star -e^{2K/3} [F (D_A W) +F^\\star (D_A W)^\\star] -3 e^KW W^\\star. \\label{FPotential} \\ee Here $K$ is the K\\\"{a}hler potential, $W(A)$ is the superpotential and $A=\\Phi\\mid_{\\theta=\\bar{\\theta}=0}$ is a complex scalar representing the lowest component of the chiral superfield $\\Phi$ (with $A$ containing the inflaton as one of its two real scalar field constituents); moreover, $D_A W = W_{,A} + K_{,A} W$ is the K\\\"{a}hler covariant derivative. Now substituting the standard solution of the auxiliary field equation \\be F=-e^{K/3}K^{,AA^\\star} (D_A W)^\\star \\label{F_StandardSolution} \\ee leads to the famous formula \\cite{Cremmer:1978hn} \\be V=e^K \\left[ K^{,AA^\\star} D_A W (D_A W)^\\star -3 W W^\\star \\right]. \\label{FamousPotential} \\ee Thus the potential is given as the difference between two manifestly positive terms (assuming the standard sign for the kinetic term, i.e. a positive K\\\"{a}hler metric $K_{,AA^\\star}$), and hence these terms are larger (and in some cases significantly larger) in magnitude than the final potential. This simple observation turns out to have important consequences in our present context: we will add the leading higher-derivative correction term to the standard kinetic term for $A$ -- in supergravity this term is of the form \\cite{Koehn:2012ar, Farakos:2012qu} \\be T (\\partial A)^2 (\\partial A^\\star)^2-2T e^{K/3} \\partial A \\cdot \\partial A^\\star FF^\\star + T e^{2K/3} (FF^\\star)^2 , \\label{X2} \\ee where $T$ has dimensions $[M^{-4}]$ and for now we take $T$ to be a small constant (when expressed in Planck units) \\footnote{Here we are suppressing spacetime indices, so for instance $(\\partial A)^2 \\equiv \\partial^\\mu A \\partial_\\mu A$.}. As one can see from the above formula, adding a higher-derivative kinetic term in supergravity brings with it corrections to the lower-order kinetic terms (second term) as well as to the potential (third term). This is the first non-trivial consequence of supersymmetry: one cannot simply build up an effective field theory term by term, since each new term modifies some of the terms previously present! The second consequence comes from estimating the leading corrections by simply plugging in the solution for $F$ that is valid in the absence of the higher-derivative terms, namely Eq. \\eqref{F_StandardSolution}, into Eq. \\eqref{X2} (with $K_{,AA^\\star=1}$): \\be T (\\partial A)^2 (\\partial A^\\star)^2-2T \\partial A \\cdot \\partial A^\\star e^{K} D_A W (D_A W)^\\star + T [e^K D_A W (D_A W)^\\star]^2 , \\label{X2-2} \\ee One would naively have expected that all corrections which the addition of the term \\eqref{X2-2} leads to are of ${\\cal O}(T)$ and that therefore, for small $T$ (say $T=1/100$), they can safely be ignored. However, as we can now see, the potential $V$ gets corrected not by a term $T V^2,$ but rather by a term proportional to $T[e^K D_A W (D_A W)^\\star]^2.$ As argued below Eq. \\eqref{FamousPotential}, this term can be significantly larger than $TV^2$ and thus the inclusion of higher-derivative {\\it kinetic} terms can, and does, lead to important modifications of the {\\it potential}. In the context of inflationary cosmology, this is evidently of importance, since the potential needs to be very flat over a significant field range. As we will see in more detail later on, such corrections have a crucial influence on the dynamics.  In the present paper, we will develop the heuristic ideas above and present a quantitative analysis of the effects of the leading higher-derivative kinetic correction term in supergravity for several representative examples. In section \\ref{structureoff} we will analyze the general structure of the theory, arguing analytically that the $T$-dependent correction to the potential will dominate over the correction to the kinetic term. We confirm this numerically by considering a number of explicit examples in section \\ref{examples}, where we study the (inflationary or not) dynamics and their dependence on $T.$ The relationship with the formalism for finding inflationary attractor solutions is discussed in section \\ref{attractor}. In section \\ref{csblowup} we study the theory for negative $T$ where there can be three real solutions to the equation of motion for $F$. We find a blowup in the speed of sound $c_s^2$ for the branch which is smoothly connected to the $T > 0$ solution. We conclude with a discussion section. Appendix A contains the relevant stability analysis for the single-field solutions that we describe in the paper. ",
        "conclusions": "In this work we have studied the effective field theory approach to cosmology in ${\\cal N}=1$ 4-dimensional supergravity, focussing on the effects of the leading higher-derivative kinetic term corrections. We would like to emphasise that from a string-theoretic perspective such terms are expected to be present on general grounds. Moreover, the cosmic microwave background data from the PLANCK satellite still leaves considerable room for the presence of such terms, as the experimental error on the associated equilateral and orthogonal non-Gaussianity parameters $f_{NL}^{equil}=-42 \\pm 75$ and $f_{NL}^{ortho}=-25 \\pm 39$ (at $2\\sigma$) \\cite{Ade:2013uln} are still quite large.  In supergravity, the leading and simplest higher-derivative terms are of the form \\cite{Koehn:2012ar} \\be T (\\partial A)^2 (\\partial A^\\star)^2-2T e^{K/3} \\partial A \\cdot \\partial A^\\star FF^\\star + T e^{2K/3} (FF^\\star)^2. \\ee The role of the auxiliary field $F$ is crucial in this context, as it leads to corrections to both the ordinary kinetic term and, most importantly, to the potential. Even in the regime where the coefficient of the square of the ordinary kinetic term is small, there can be dramatic changes occurring in the potential, which drastically modify any possible inflationary dynamics. We have restricted our analysis to the case where $T$ is a constant. The general formalism \\cite{Koehn:2012ar} would allow $T$ to also contain derivatives of the fields, e.g. $T$ could contain a factor $(-\\partial A \\cdot \\partial A^\\star)^n \\supset X^n$ with $n>0.$ As is evident from the expression above, such a term would then also yield corrections to the $X^{n+1}$ and $X^n$ terms -- however, it would not modify the form of the potential term (except of course for the change it induces in the solution for $F$). In an inflationary context, the non-derivative $T$ corrections are thus the most important ones.  For $T \\geq 0$, we gave the general solution for $F$. We showed that the corrections can be large even when $T$ is small -- in particular, the corrections cannot be ignored when \\begin{eqnarray} 32 T e^K K^{A A^\\star} |D_A W|^2 &\\gtrsim&1. \\end{eqnarray} We then compared the $T$-dependent kinetic and potential term contributions, finding  that the T-dependent contribution to the potential will be more significant than the T-dependent contribution to the kinetic part of the action. We have also verified that the attractor formalism developed in \\cite{Franche:2009gk} remains valid for these supergravity $P(X , \\phi)$ actions. We considered several explicit examples in detail, confirming that for these even very small increases in $T$ can result in significant changes to the potential and to the inflationary dynamics.  Our findings suggest that it is essential to take into account the full supergravity description of higher-derivative terms, rather than assuming that a small parameter expansion is possible in which the only relevant terms are $X^n$ etc. If a potential is already unstable, the higher-derivative $T$-dependent corrections tend to make it even more unstable. In the inflationary context, the higher-derivative terms play a more welcome role when the potential curves upwards too steeply (i.e. when $V_{,\\phi\\phi}>0$), since in such a situation they can lower and flatten the potential. In addition the higher-derivative kinetic term slows the field down, as expected in non-canonical inflationary scenarios. On the other hand the effect of higher-derivative corrections on the potential may be useful in ekpyrotic models also, where a steep negative potential is required \\cite{Khoury:2001wf, Lehners:2008vx} -- we leave such an investigation for future work.  For $T <0$, we gave general expressions for the three solutions to the cubic for $F$ in section \\ref{Tneg}, and determined which is real for the three relevant $T <0$ regions defined by the relative values of $T$ and $X$. However,  the system seems to be blocked from the region with three real solutions for $T <0$ by a blow-up of $c_s^2$, the speed of propagation of perturbations, which should have $c_s=1$ as an upper limit.  This study represents an important step forward in the understanding of how higher-order kinetic terms should be embedded in a consistent (cosmological) supergravity context. It will be interesting to look for set-ups in which these effects lead to a useful model of inflationary or ekpyrotic cosmology, consistent with the available data but not suffering from excessive fine tuning. We hope to be able to return to this question in the future."
    },
    "1402/1402.0533_arXiv.txt": {
        "abstract": "We review models of chondrite component transport in the gaseous protoplanetary disk. Refractory inclusions were likely transported by turbulent diffusion and possible early disk expansion, and required low turbulence for their subsequent preservation in the disk, possibly in a dead zone. Chondrules were produced locally but did not necessarily accrete shortly after formation. Water may have been enhanced in the inner disk because of inward drift of solids from further out, but likely not by more than a factor of a few. Incomplete condensation in chondrites may be due to slow reaction kinetics during temperature decrease. While carbonaceous chondrite compositions might be reproduced in a ``two-component'' picture (Anders 1964), such components would not correspond to simple petrographic constituents, although part of the refractory element fractionations in chondrites may be due to the inward drift of refractory inclusions. Overall, considerations of chondrite component transport alone favor an earlier formation for carbonaceous chondrites relative to their noncarbonaceous counterparts, but independent objections have yet to be resolved. ",
        "introduction": "With the accelerating pace of exoplanet detections% , the protoplanetary disk phase of stellar evolution enjoys considerable interest. Thanks to increasing computational power, theorists can test mechanisms for disk transport \\citep{Turneretal2014} and planet formation \\citep{YoudinKenyon2013}. Observations of present-day protoplanetary disks \\citep{WilliamsCieza2011} probe the disk mass, size, structure, chemical species and solids \\citep{Nattaetal2007}. However, even with the Atacama Large Millimeter/Submillimeter Array% , it will remain challenging to resolve scales below a few AUs and probe the optically thick midplane of the inner disks where planet formation should occur.  To understand the evolution of solids in disks, we must turn our attention to constraints provided closer to us by our own solar system in the form of primitive meteorites, or \\textit{chondrites}. Indeed, chondrites date back to the protoplanetary disk phase of the solar system, 4.57 Ga ago, and with more than 40,000 specimens classified to date, not to mention samples returned from comet Wild 2 \\citep[e.g.][]{Zolenskyetal2008Wild2} or asteroid % Itokawa \\citep{Nakamuraetal2011}, they offer a considerable wealth of petrographic, chemical, and isotopic data at all examination scales.  Yet chondrites arrive in our laboratories without geological context. While orbit determinations consistently assign their parent bodies to the asteroid main belt---with the exception of micrometeorites \\citep[e.g.][]{EngrandMaurette1998} and perhaps some carbonaceous chondrites \\citep{Gounelleetal2008} possibly derived from further out---, exactly where and when they originally accreted is largely unknown. Still, chondrites exhibit considerable compositional variations, and space and time were obviously important dimensions behind them. In fact, each individual chondrite is a mixture of components (chondrules, refractory inclusions, etc.) formed in different locations, epochs and environments in the disk \\citep{BrearleyJones1998, Krotetal2009}. This is evidence for considerable transport in the disk. In order to place the meteoritical record in context, the relevant transport processes have to be understood, and as such meteorites are sensors of the dynamics of protoplanetary disks.  Our purpose here is to review transport mechanisms of chondrite components before accretion. An earlier review on particle-gas dynamics was given by \\citet{CuzziWeidenschilling2006} and \\citet{Boss2012} reviewed transport and mixing from the perspective of isotopic heterogeneity. The formation \\textit{per se} of chondrite components is essentially beyond our scope but the reader may be referred to recent reviews by \\citet{Krotetal2009} and \\citet{Aleon2010}. \\citet{Wood2005} and \\citet{Chambers2006} proposed syntheses on the origin of chondrite types from cosmochemical and astrophysical viewpoints, respectively. Here, the discussion will be organized around meteoritical constraints as follows: In section \\ref{Background}, we provide background on chondrites and the basic physics of the protoplanetary disk before embarking in section \\ref{Transport} on an examination of transport constraints from specific chondrite components. We then review the interpretation of fractionation trends exhibited by chondrites as wholes (section \\ref{Fractionation}) in light of which we will discuss the chronological and/or spatial ordering of chondrite groups (section \\ref{Chondrites in space and time}). ",
        "conclusions": ""
    },
    "1402/1402.3530_arXiv.txt": {
        "abstract": "Using archival data from the \\mbox{H\\,{\\sc i}} Parkes All Sky Survey (HIPASS) we have searched for 21\\,cm line absorption in 204 nearby radio and star-forming galaxies with continuum flux densities greater than $S_{1.4} \\approx 250$\\,mJy within the redshift range $0 < cz < 12\\,000\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$. By applying a detection method based on Bayesian model comparison, we successfully detect and model absorption against the radio-loud nuclei of four galaxies, of which the Seyfert~2 galaxy 2MASX\\,J130804201-2422581 was previously unknown. All four detections were achieved against compact radio sources, which include three active galactic nuclei (AGNs) and a nuclear starburst, exhibiting high dust and molecular gas content.  Our results are consistent with the detection rate achieved by the recent ALFALFA (Arecibo Legacy Fast Arecibo L-band Feed Array) \\mbox{H\\,{\\sc i}} absorption pilot survey by Darling et al. and we predict that the full ALFALFA survey should yield more than three to four times as many detections as we have achieved here. Furthermore, we predict that future all-sky surveys on the Square Kilometre Array precursor telescopes will be able to detect such strong absorption systems associated with type 2 AGNs at much higher redshifts, providing potential targets for detection of H$_{2}$O megamaser emission at cosmological redshifts. ",
        "introduction": "\\label{section:introduction} Atomic hydrogen (\\mbox{H\\,{\\sc i}}) gas, traced by the 21\\,cm line, is a powerful probe of the mass distribution within galaxies and the available fuel for future star formation. However, the strength of the 21\\,cm emission-line decreases rapidly with increasing redshift as a function of the inverse square of the luminosity distance. In individual galaxies the most distant detectable 21\\,cm emission lines are at $z\\sim0.2$ \\citep[e.g.][]{Catinella:2008, Verheijen:2010, Freudling:2011}, while statistical detections using spectral stacking have reached $z\\approx0.37$ \\citep{Lah:2009}. At significantly higher redshifts, we can instead detect the 21\\,cm line in individual galaxies through the absorption of continuum flux towards a background radio source. In principle, \\mbox{H\\,{\\sc i}} absorption can be observed up to cosmological redshifts, where the ionosphere begins to corrupt the signal, yet such observations are ultimately limited by the sample of known high-redshift radio sources and the availability of suitable instrumentation. The highest redshifts achieved by observations of 21\\,cm absorbers include the radio galaxy B2\\,0902+34 ($z = 3.397$; \\citealt{Uson:1991}) and the intervening system towards the quasar PKS\\,0201+113 ($z = 3.387$; \\citealt{Kanekar:2007}). The detection limit for a survey of 21\\,cm line absorption is independent of redshift, and depends only on the availability of bright background continuum sources against which the line can be detected. For any such sight line, absorption is particularly sensitive to high column densities ($N_\\mathrm{HI} \\gtrsim 10^{20}\\,\\mathrm{cm}^{-2}$) of cold ($T_\\mathrm{spin} \\lesssim 100$\\,K) foreground \\mbox{H\\,{\\sc i}} gas that obscures a large fraction of the background radio source.  In the local Universe, at least 10\\,\\,per\\,cent of extragalactic radio sources that have been searched exhibit an associated 21\\,cm absorption line at or near the optical redshift, indicating that neutral gas is present within the host galaxy \\citep[e.g.][]{Morganti:2001, Vermeulen:2003, Allison:2012a}. High signal-to-noise ratio (S/N) absorption lines, typically associated with powerful radio galaxies, often exhibit broad wings, which can indicate the presence of fast jet-driven outflows of \\mbox{H\\,{\\sc i}} gas (with velocities over 1000 km\\,s$^{-1}$ and outflow rates of several tens of M$_\\odot$\\,yr$^{-1}$; e.g. \\citealt{Morganti:2005b,Mahony:2013,Morganti:2013}). These high-velocity outflows may have a profound effect on the star formation and subsequent evolution of the host galaxy. Furthermore, there is evidence to suggest that in some cases broad absorption components can arise in circumnuclear gas distributed as a disc or torus \\citep[e.g.][]{Struve:2010a, Morganti:2011}. Such observations are incredibly useful for directly studying the interaction between the radio-loud nucleus and the neutral gas in the interstellar medium. However, at present these surveys are limited to targeted sampling of the radio source population, typically focusing on those that are compact, and by doing so can introduce biases \\citep[see e.g.][]{Curran:2010}.  Here we present the results of a search for 21\\,cm absorption in nearby radio and star-forming galaxies from the \\mbox{H\\,{\\sc i}} Parkes All-Sky Survey (HIPASS; \\citealt{Barnes:2001}). While the data are considerably less sensitive than current targeted observations of radio sources (with an effective integration time of 7.5\\,min per individual pointing and typical rms noise of $\\sim$13\\,mJy\\,beam$^{-1}$ per 13\\,km\\,s$^{-1}$ channel separation), the large volume covered by HIPASS (the whole sky south of $\\delta = +25^\\circ$ and $z\\lesssim0.042$) allows identification of the strongest associated \\mbox{H\\,{\\sc i}} absorption-line systems in the local universe in an unbiased way. This enables us to study some of the most extreme and potentially interesting systems, as well as testing line-finding techniques \\citep[e.g.][]{Allison:2012b} that can be used in planning future, more sensitive, large-area surveys with the Square Kilometre Array (SKA) pathfinder and precursor telescopes.  \\cite{Darling:2011} recently published the results of a pilot survey for \\mbox{H\\,{\\sc i}} 21\\,cm absorption in the Arecibo Legacy Fast Arecibo L-band Feed Array (ALFALFA) survey.  This was the first genuinely blind search for absorption within a large-area radio survey, and covered 517\\,deg$^2$ of sky in the redshift range $z<0.058$. No intervening lines were seen, but one previously known associated line was re-detected in the interacting luminous infrared galaxy UGC\\,6081 \\citep{Bothun:1983, Williams:1983}. The HIPASS search presented here can be considered complementary to that survey, since it covers a much larger area of sky (by approximately a factor of 50) with similar redshift coverage, but has lower sensitivity (by approximately a factor of 6). However, due to the presence of strong baseline ripple, we have limited our search to the detection of \\mbox{H\\,{\\sc i}} absorption within the host galaxies of the radio sources themselves. Spectral baseline ripples are a common problem for single-dish observations of the 21\\,cm line \\citep[e.g.][]{Briggs:1997} and HIPASS spectra towards bright continuum sources are particularly affected, where standing waves are generated between the primary dish and receiver cabin \\citep{Barnes:2001,Barnes:2005}. By using the known systemic redshift of the galaxy as a prior, we can attempt to distinguish the absorption-line from the strong baseline ripple. We intend in future work to revisit the HIPASS data with improved analysis and perform an extended search of intervening \\mbox{H\\,{\\sc i}} absorption within the full volume.  Throughout this paper we adopt a flat $\\Lambda$ cold dark matter cosmology with $H_{0}$ = 70\\,km\\,s$^{-1}$, $\\Omega_\\mathrm{M}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7. Radial velocities and redshifts have been corrected for the solar barycentric standard-of-rest frame. ",
        "conclusions": "\\begin{figure*} \\centering \\includegraphics[width = 1.0\\textwidth]{cabb_spectra.pdf} \\caption{CABB 21\\,cm spectra showing both continuum emission and \\mbox{H\\,{\\sc i}} absorption at the positions of radio sources hosted by 2MASX\\,J13084201-2422581 and NGC\\,5793. The arrow and horizontal line indicate the mean and uncertainty in the systemic redshift. The vertical dashed lines enclose those spectral channels that were flagged either due to persistent strong RFI or Galactic 21\\,cm signal. The emission spike that is apparent in both spectra, within a single channel at the band centre (1406\\,MHz, $cz \\approx 3100\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$), arises due to self-generated interference within the telescope (see \\citealt{Wilson:2011}).}\\label{figure:cabb_spectra} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width = 1.0\\textwidth]{cabb_images.pdf} \\caption{1.4\\,GHz ATCA images of 2MASX\\,J13084201-2422581 and NGC\\,5793, where the blue contours denote the 5, 10, 20, 50\\,per\\,cent levels of the peak continuum flux density. The FWHM of the synthesized beam is shown in the bottom-right corner. The grey-scale images represent optical $B_{j}$-band photometry from the SuperCosmos Sky Survey, using the UK Schmidt and Palomar Oschin Schmidt telescopes \\citep{Hambly:2001}.} \\label{figure:image_overlays_cabb} \\end{figure*}  \\begin{figure*} \\vspace{6pt} \\centering \\includegraphics[width = 1.0\\textwidth]{spectral_fits.pdf} \\caption{The best-fitting models of HIPASS and CABB spectra in which we have detected associated \\mbox{H\\,{\\sc i}} absorption. The radial velocity axis is given relative to the rest frame defined by the systemic redshift (vertical dotted line). For the CABB spectra, the grey line represents the data after we have subtracted our best-fitting continuum component, whereas for the HIPASS spectra the grey line represents the original baseline-subtracted data. The solid black line represents the best-fitting spectral line model and the dashed green lines the individual Gaussian components (top panel). The solid red line represents the best-fitting residual (bottom panel).}\\label{figure:spectral_fits} \\end{figure*}  \\subsection{Individual sources}\\label{section:individual_sources}  \\subsubsection{2MASX\\,J13084201-2422581}  We detect a previously unknown 21\\,cm absorption line against the compact flat-spectrum radio source at the centre of the Seyfert~2 galaxy 2MASX\\,J13084201-2422581. We show in Fig.\\,\\ref{figure:spectral_fits} the best fitting models to both the HIPASS and CABB spectra. By comparing the marginal likelihoods for increasingly complex models, we find that the HIPASS data warrant only a single-component Gaussian model, while a four-component model is favoured by the CABB data. The similarity in our estimates of the peak depth and rest effective width for each spectrum implies that all of the \\mbox{H\\,{\\sc i}} absorption detected within the 15.5\\,arcmin HIPASS beam arises from a region of angular size smaller than the ATCA synthesized beam, which at $cz = 4257$\\,km\\,s$^{-1}$ equates to a projected physical size smaller than $6 \\times 3$\\,kpc. This result is consistent with the compact morphology of the radio source at 1.4\\,GHz, evident from the NVSS and ATCA images, and the absence of other nearby strong radio sources within the HIPASS beam (see Fig.\\,\\ref{figure:image_overlays_detections}).  The redshift of the peak absorption is consistent with the systemic redshift of the host galaxy, implying that the bulk of the cold $\\mbox{H\\,{\\sc i}}$ gas is not rapidly infalling or outflowing with respect to ionized gas in the nucleus. The stellar component of the host galaxy exhibits an edge-on irregular spiral morphology at near-infrared and optical wavelengths \\citep{Jarrett:2000, Hambly:2001}, which, with the Seyfert~2 classification of the AGN, suggests that the bulk of the absorbing gas may arise within an obscuring disc of \\mbox{H\\,{\\sc i}} gas. While we cannot spatially resolve the background radio-jet structure with our ATCA observations, and hence strongly constrain the spatial distribution and kinematics of the \\mbox{H\\,{\\sc i}} gas, we note that the shape and width of the absorption-line profile are very similar to those observed in other Seyfert galaxies (e.g. \\citealt{Dickey:1982, Dickey:1986, Gallimore:1999}). Work by \\citet{Gallimore:1999} showed that these systems are well modelled by sub-kpc discs of \\mbox{H\\,{\\sc i}} gas that are typically aligned with the outer stellar disc. Deviations in the regular shape of the main profile (components 2, 3 and 4 in Fig.\\,\\ref{figure:spectral_fits}) are likely generated by a combination of unresolved spatial variations in the optical depth of the gas, the complex geometries of the absorber-radio source system and radial streaming of the gas with respect to the source. The separate broad and shallow blueshifted component at $v\\sim200\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ indicates that gas might be caught in a jet-driven outflow on sub-kpc scales, but this interpretation remains tentative until the absorption can be spatially resolved.  Our best estimates of the peak and integrated 21\\,cm optical depths from the CABB data are $0.28\\pm0.01$ and $23.3\\pm0.5$\\,km\\,s$^{-1}$, respectively (assuming that $f=1$). However, without further knowledge of the relative size and geometry of the absorbing gas with respect to the continuum source, as well as the spin temperature of gas, it is very difficult to obtain an accurate measurement of the column density from Equation\\,\\ref{equation:column_density}. \\cite{Gallimore:1999} showed that for a dense AGN-irradiated gas cloud in the narrow-line region of a Seyfert galaxy, the 21\\,cm spin temperature is likely to be collisionally dominated with typical values of $T_\\mathrm{spin} = 100$\\,K. However, if the sightline to the continuum source intercepts the warmer atomic medium, then the spin temperature may be much higher. For example, 21\\,cm observations of intervening damped Lyman $\\alpha$ absorbers ($N_\\mathrm{HI} > 2 \\times 10^{20}\\,\\mathrm{cm}^{-2}$) show that $T_\\mathrm{spin}/f$ can be greater than $1000\\,\\mathrm{K}$ (see \\citealt{Curran:2012} and references therein). Furthermore, if the \\mbox{H\\,{\\sc i}} gas in this Seyfert~2 galaxy is distributed as a disc on scales less than 100\\,pc, then it would be unlikely that all of the source structure would be uniformly obscured by the absorbing gas, and so in this case we would expect the covering factor to be less than unity. Therefore, given the possible values of $T_\\mathrm{spin}$ and $f$, we can only estimate a lower limit to the \\mbox{H\\,{\\sc i}} column density of $N_\\mathrm{HI} = 42.5\\pm0.8 \\times 10^{20}~(T_\\mathrm{spin}/100\\,\\mathrm{K})~\\mathrm{cm}^{-2}$. The stellar disc evident in the 2MASS $K_{s}$-band photometry for this galaxy has a major--minor axis ratio of 0.380 \\citep{Jarrett:2000}, which we convert into an inclination angle of $i = 74\\degr$ \\citep{Tully:1977, Aaronson:1980}. Assuming that the \\mbox{H\\,{\\sc i}} gas is coplanar with the stellar component, our estimate of the column density is consistent with the inclination angle relationship measured in other Seyferts and active galaxies by \\cite{Dickey:1982,Dickey:1986} and \\cite{Gallimore:1999}.  \\begin{table} \\begin{threeparttable} \\caption{Best-fitting parameters from fitting multiple Gaussian components to the HIPASS and CABB 21\\,cm spectra. $n$ is the component number, corresponding to the label given in Fig.\\,\\ref{figure:spectral_fits}; $cz$ is the component redshift; $\\Delta{v_{\\rm FWHM}}$ is the velocity FWHM and $\\Delta{S}$ is the depth.}\\label{table:parameters} \\setlength\\tabcolsep{2pt} \\setlength\\tabcolsep{1pt} \\begin{tabular}{@{}llrrrr@{}} \\hline \\multicolumn{1}{l}{Name} & \\multicolumn{1}{l}{Data} & \\multicolumn{1}{c}{$n$} & \\multicolumn{1}{c}{$cz$} & \\multicolumn{1}{c}{$\\Delta{v_{\\rm FWHM}}$} & \\multicolumn{1}{c}{$\\Delta{S}$} \\\\ & & & \\multicolumn{1}{c}{(km\\,s$^{-1}$)} & \\multicolumn{1}{c}{(km\\,s$^{-1}$)} & \\multicolumn{1}{c}{(mJy)} \\\\ \\hline 2MASX\\,J13084201-2422581 & HIPASS & 1 & 4224 & 66 & 102 \\\\ & CABB & 1 & 4050 & 127 & 6 \\\\ & & 2 & 4170 & 12 & 53 \\\\ & & 3 & 4219 & 65 & 114 \\\\ & & 4 & 4243 & 18 & 36 \\\\ Centaurus\\,A & HIPASS & 1 & 548 & 3 & 2686 \\\\ & & 2 & 564 & 21 & 241 \\\\ NGC\\,5793 & HIPASS & 1 & 3484 & 63 & 702 \\\\ & & 2 & 3486 & 380 & 36 \\\\ & & 3 & 3526 & 39 & 544 \\\\ & & 4 & 3569 & 25 & 325 \\\\ & CABB & 1 & 3474 & 15 & 143 \\\\ & & 2 & 3496 & 83 & 848 \\\\ & & 3 & 3525 & 35 & 310 \\\\ & & 4 & 3570 & 24 & 298 \\\\ Arp\\,220 & HIPASS & 1 & 5482 & 348 & 42 \\\\ \\hline \\end{tabular} \\end{threeparttable} \\end{table}  \\begin{figure} \\centering \\includegraphics[width = 1.0\\columnwidth]{cenA_compare.pdf} \\caption{HIPASS spectra towards Centaurus\\,A, extracted at 15.5\\,arcmin separations along the observed jet-axis, at RA (J2000) = $13^{\\rm h}26^{\\rm m}27\\fs72$, Dec. (J2000) = $-42\\degr50\\arcmin09\\farcs8$ (north), RA (J2000) = $13^{\\rm h}25^{\\rm m}27\\fs76$, Dec. (J2000) = $-43\\degr01\\arcmin07\\farcs4$ (centre) and RA (J2000) = $13^{\\rm h}24^{\\rm m}27\\fs80$, Dec. (J2000) = $-43\\degr12\\arcmin05\\farcs0$ (south). The radial velocity axis is given relative to the rest frame defined by the systemic redshift (vertical dashed line).}\\label{figure:cenA_compare} \\end{figure}  \\subsubsection{Centaurus\\,A}  Centaurus\\,A (NGC\\,5128) is by far the closest early-type radio galaxy to the Milky Way, which at a distance of only 3.8\\,Mpc \\citep{Harris:2010} has been imaged in detail at multiple wavelengths (see \\citealt{Israel:1998} and references therein). In the HIPASS spectrum, extracted from the core of the radio source, we re-detect the \\mbox{H\\,{\\sc i}} absorption first discovered by \\cite{Roberts:1970} and studied extensively since \\citep[e.g.][]{Whiteoak:1971, VanderHulst:1983, Sarma:2002, Morganti:2008, Struve:2010c}. Due to the proximity and radio power ($P_{1.4} \\approx 5 \\times 10^{24}$\\,W\\,Hz$^{-1}$) of this source, the HIPASS spectrum is strongly contaminated by the spectral baseline ripple. Despite this, we recover a two-component Gaussian model of the line (Fig.\\,\\ref{figure:spectral_fits}c), with an effective width of $7.4_{-2.3}^{+3.6}$\\,km\\,s$^{-1}$ and a peak depth of $1690_{-550}^{+770}$\\,mJy. The poor constraints on these parameter estimates, compared with those for the other detections, are the result of low spectral sampling across the line.  The structure of the line profile, with a deep narrow component at the systemic redshift and an broadened component towards higher redshifts, is consistent with the structure seen at similar spectral and spatial resolution by \\cite{Roberts:1970} and \\cite{Whiteoak:1971}. The deeper narrow component is thought to arise in absorption from \\mbox{H\\,{\\sc i}} gas in a rotating disc that is coplanar with the prominent warped dust lane, while some of the redshifted absorption is consistent with infalling clouds towards the nucleus \\citep{VanderHulst:1983, Sarma:2002}. Higher spatial resolution and more sensitive observations by \\cite{Morganti:2008} and \\cite{Struve:2010c} revealed blueshifted absorption towards the nucleus, potentially indicating the presence of a circumnuclear disc of \\mbox{H\\,{\\sc i}} on sub-100\\,pc scales.  The 1.4\\,GHz emission from Centaurus\\,A is moderately extended with respect to the HIPASS beam, and so in Fig.\\,\\ref{figure:cenA_compare} we show spectra extracted at three positions along the observed jet axis, separated by intervals of 15.5\\,arcmin. While evidently contaminated by residual spectral baseline signal, and subject to adjacent signal entering from the beam sidelobes, we tentatively see more absorption towards the southern end of the jet axis. This is consistent with the orientation of the \\mbox{H\\,{\\sc i}} disc against the southern radio lobe, seen at higher spatial resolution \\citep[e.g.][]{Struve:2010c}.  \\subsubsection{NGC\\,5793}  The very deep 21\\,cm absorption seen towards the compact and radio-luminous nucleus in this edge-on disc Seyfert~2 galaxy was first detected by \\cite{Jenkins:1983} using the 64\\,m Parkes Radio Telescope, and has since been studied at higher spatial resolution by \\cite{Gardner:1986} using the Very Large Array (VLA), and by \\cite{Gardner:1992} and \\cite{Pihlstrom:2000} using very long baseline interferometry (VLBI). The absorption-line profile seen in the HIPASS spectrum \\citep{Koribalski:2004} is consistent with that observed by \\cite{Jenkins:1983}. The weaker emission-line feature seen in both spectra at $cz \\approx 2860$\\,km\\,s$^{-1}$ is attributed by \\cite{Koribalski:2012} to \\mbox{H\\,{\\sc i}} gas in the neighbouring dwarf irregular galaxy 6dF\\,J1459410-164235 (to the east), and not the E0 galaxy NGC\\,5796 (to the north), which is thought to be relatively \\mbox{H\\,{\\sc i}} poor.  \\begin{figure*} \\vspace{6pt} \\centering \\includegraphics[width = 1.0\\textwidth]{luminosity.pdf} \\caption{Left: the 1.4\\,GHz radio power versus the systemic redshift for all 204 galaxies in our sample. Right: the 1.4\\,GHz radio power versus the far-infrared luminosity for those 86 galaxies with far-infrared 60 or 100\\,$\\mathrm{\\mu}$m flux densities from the \\emph{IRAS} survey. The blue squares represent those galaxies with \\mbox{H\\,{\\sc i}} emission catalogued in HIPASS and the red circles denote those in which we have detected absorption. The dashed line shows the radio power versus far-infrared luminosity relationship measured by \\citet{Mauch:2007} for star-forming galaxies. The dotted lines indicate a one-decade deviation from this relationship.}\\label{figure:luminosity} \\end{figure*}  The spatially unresolved absorption lines in both our HIPASS and CABB spectra clearly exhibit some velocity structure, which we successfully model using a four-component Gaussian model (see Fig.\\,\\ref{figure:spectral_fits}b). VLBI observations by \\cite{Pihlstrom:2000} demonstrated that this structure results from the superposition of individual \\mbox{H\\,{\\sc i}} components seen against two continuum sources that are only resolved on angular scales smaller than 10\\,mas. They suggest that the broadest feature likely arises in a nearly edge-on disc of \\mbox{H\\,{\\sc i}} gas ($i \\approx 73\\,\\degr$), and occurs on scales of 50--100\\,pc from the AGN, consistent with that seen in other Seyfert~2s \\citep{Gallimore:1999}, while the other features are signatures of individual \\mbox{H\\,{\\sc i}} clouds that are either interior or exterior to this disc. By fitting a four component Gaussian model to the CABB spectrum, we estimate that the peak and integrated optical depth are $2.32\\pm0.03$ and $157.6\\pm0.5$\\,km\\,s$^{-1}$, respectively, giving a total \\mbox{H\\,{\\sc i}} column density of $N_\\mathrm{HI} = 287.2\\pm1.0 \\times 10^{20} (T_\\mathrm{spin}/100\\,\\mathrm{K})$\\,cm$^{-2}$. This is consistent with the total column density measured by \\cite{Pihlstrom:2000}, averaged across the resolved continuum components, of $N_\\mathrm{HI} \\approx 3.5 \\times 10^{22} (T_\\mathrm{spin}/100\\,\\mathrm{K})$\\,cm$^{-2}$.  \\cite{Koribalski:2012} tentatively identified a previously undetected broad absorption feature in the HIPASS spectrum, with a width of $680$\\,km\\,s$^{-1}$ and centred on the systemic redshift. Using Bayesian model comparison, we confirm that this feature is statistically significant above the noise (component 1 in Fig.\\,\\ref{figure:spectral_fits}b and Table\\,\\ref{table:parameters}); however, it is not clear if this feature is distinguishable from other residual baseline features seen in the spectrum. Furthermore, we do not re-detect this broad component in the CABB spectrum (with an estimated per-channel noise of $\\sigma_{\\rm chan} = 3.9$\\,mJy), even though there is strong consistency between the other absorption components seen in both spectra. It is plausible that this feature could have arisen towards a confused source within the HIPASS beam, from \\mbox{H\\,{\\sc i}} gas that is at a similar redshift to NGC\\,5793. There are two other sources within the HIPASS beam that have sufficient flux densities ($S_{1.4} \\gtrsim 35$\\,mJy) in the NVSS catalogue to produce such an absorption: NGC\\,5796 ($S_{1.4} = 109$\\,mJy, $cz_{\\rm sys} = 2971$\\,km\\,s$^{-1}$; \\citealt{Wegner:2003}) and MRC\\,1456-165 ($S_{1.4} = 379$\\,mJy).  However, spectra extracted from the CABB data at the centroid positions of both sources show no evidence of the broad absorption seen in the HIPASS spectrum. We therefore conclude that this feature is likely an artefact and the result of residual spectral baseline ripple in the HIPASS spectrum.  \\subsubsection{Arp\\,220}  The broad absorption-line associated with this prototypical ultraluminous infrared galaxy (\\citealt{Sanders:2003}) was originally detected by \\cite{Mirabel:1982}, using the 300\\,m Arecibo Telescope, and has since been re-observed and studied multiple times (e.g. \\citealt{Dickey:1986, Baan:1987, Garwood:1987, Baan:1995, Hibbard:2000, Mundell:2001}). We find that the line detected in the HIPASS spectrum requires only a single-component Gaussian model, with an effective width ($\\Delta{v_{\\rm eff}} = 375_{-41}^{+46}$\\,km\\,s$^{-1}$) and peak depth ($\\Delta{S_{\\rm peak}} = 42.5_{-4.5}^{+4.7}$\\,mJy) that are consistent with previous single-dish observations (e.g \\citealt{Mirabel:1982, Garwood:1987}). However, the lower S/N and spectral resolution of the HIPASS spectrum means that we do not find as much structure in the line as seen in these other single-dish observations.  The absorption arises from gas towards a compact radio-loud nucleus that consists of two distinct components (e.g. \\citealt{Baan:1987, Norris:1988, Baan:1995}), which are thought to be the nuclei of two gas-rich progenitor galaxies in an advanced stage of merging. At 1.4\\,GHz, they are only resolved on angular scales smaller than $\\sim200$\\,mas \\citep[][]{Mundell:2001} and are therefore not resolved by HIPASS. Mundell et al. carried out a high spatial resolution study of the 21\\,cm line on sub-arcsec scales, using the Multi-Element Radio-Linked Interferometer Network array, and showed that the bulk of the absorption is likely associated with two counterrotating discs of \\mbox{H\\,{\\sc i}} gas centred on each of the nuclei, consistent with observations of emission from the CO gas content \\citep{Sakamoto:1999}. The broad width of the absorption line seen in the HIPASS spectrum is consistent with the superposition of these rotating components and the bridge of \\mbox{H\\,{\\sc i}} gas connecting the two nuclei.  \\begin{figure} \\vspace{6pt} \\centering \\includegraphics[width = 1.0\\columnwidth]{luminosity_class.pdf} \\caption{The distribution of 1.4\\,GHz radio powers for our sample, where we have classified galaxies as either star forming or AGN dominated (based on the radio and far-infrared luminosity relationship), and which have radio morphologies as classified by \\citet{VanVelzen:2012}. The hatched bars represent those galaxies with catalogued 21\\,cm line emission in HIPASS (\\citealt{Koribalski:2004, Meyer:2004, Wong:2006}).}\\label{figure:luminosity_class} \\end{figure}  \\begin{figure} \\vspace{6pt} \\centering \\includegraphics[width = 1.0\\columnwidth]{redshift_class.pdf} \\caption{As Fig.\\,\\ref{figure:luminosity_class}, but for the distribution of redshifts.}\\label{figure:redshift_class} \\end{figure}  \\subsection{Comparison with the 2\\,Jy sample}\\label{section:2jy_sample}  \\cite{Morganti:2001} used the ATCA, the VLA and the Westerbork Synthesis Radio Telescope to search for \\mbox{H\\,{\\sc i}} absorption in 23 radio galaxies (at $z<0.22$ and $\\delta < +10\\degr$) selected from the 2\\,Jy sample \\citep{Wall:1985}. We can use this relatively homogeneous set of observations to determine whether our non-detections in HIPASS are consistent with what we would expect from existing detections of absorption. In five of these radio galaxies, Morganti et al. detected \\mbox{H\\,{\\sc i}} absorption, of which NGC\\,5090 ($cz = 3421\\pm21$\\,km\\,s$^{-1}$) and 3C\\,353 ($cz = 9120\\pm59$\\,km\\,s$^{-1}$) are within the volume surveyed by HIPASS. The 21\\,cm spectra of these two galaxies exhibit peak absorption of 8 and 10\\,mJy, with FWHMs of approximately 100 and 200\\,km\\,s$^{-1}$, respectively. Given the noise and baseline ripple confusion in the HIPASS spectra, our non-detection of \\mbox{H\\,{\\sc i}} absorption in these radio galaxies is consistent with the expected strength of these lines.  \\subsection{The HIPASS detection rate}  We obtain detection rates for associated absorption in HIPASS of 2.0\\,per\\,cent (4/204) for the total sample, 4.4\\,per\\,cent (4/90) for Sample 1 and 1.6\\,per\\,cent (3/189) for Sample 2. While such a small number of detections does not allow us to draw strong conclusions about the population, we can attempt to understand these rates in the context of the HIPASS survey parameters and the properties of individual galaxies in the sample.  In Fig.\\,\\ref{figure:luminosity}, we show the 1.4\\,GHz radio power versus systemic redshift (for all 204 galaxies in our sample), and the far-infrared luminosity for those 86 galaxies that either have a detection at 60 or 100\\,$\\mathrm{\\mu}$m in the \\emph{Infrared Astronomical Satellite} (\\emph{IRAS}) Faint Source, Point Source and Galaxy Catalogues \\citep{Beichman:1988, Rice:1988, Knapp:1989, Moshir:1992, Sanders:2003}. The radio power is estimated using the larger of either the 843\\,MHz SUMSS (assuming a spectral index of $-0.6$) or 1.4\\,GHz NVSS total flux densities, and thereby accounting for components that might be present in SUMSS but missing in the NVSS images.  We calculate the far-infrared luminosity using an estimate of the flux density ($S_\\mathrm{FIR}$) between 42.5 and 122.5\\,$\\mathrm{\\mu}$m, which is given by \\citep{Helou:1985} \\begin{equation} S_\\mathrm{FIR} = 1.26 \\times 10^{-14}\\,(2.58\\,S_\\mathrm{60\\mu m} + S_\\mathrm{100\\mu m})\\,\\mathrm{W\\,m^{-2}}, \\end{equation} where $S_\\mathrm{60\\mu m}$ and $S_\\mathrm{100\\mu m}$ are the 60 and 100\\,$\\mu$m flux densities in units of Jy. For those galaxies where measurements of only $S_\\mathrm{60\\mu m}$ or $S_\\mathrm{100\\mu m}$ are available, we use $\\langle\\log_{10}(S_\\mathrm{100\\mu m}/S_\\mathrm{60\\mu m})\\rangle = 0.3$, which is the average calculated from the \\emph{IRAS} Bright Galaxy sample by \\cite{Soifer:1989}. Galaxies that are identified as star forming exhibit a strong correlation between their radio and far-infrared luminosities \\citep[e.g][]{Helou:1985, Devereux:1989, Condon:1991}. For a large sample of spectroscopically identified star-forming galaxies, \\cite{Mauch:2007} measured this relationship to be \\begin{equation} \\log_{10}(P_{1.4}) = (1.06 \\pm 0.01) \\log_{10}(L_{\\mathrm{FIR}}) + (11.1 \\pm 0.1), \\end{equation} with a maximum deviation in radio power of approximately one decade. We use this to identify galaxies in our sample that are star-forming, classifying AGN dominated galaxies as those that do not follow this relationship or do not have detections in both the 60 and 100\\,$\\mathrm{\\mu}$m bands.  Based on the $P_{1.4}$--$L_\\mathrm{FIR}$ relation, we estimate that 20\\,per\\,cent (47/204) of our galaxies are star forming, which is consistent with the 42 predicted using the local radio luminosity function of \\cite{Mauch:2007}. Considering the relative radio and near-infrared morphologies of these star-forming galaxies, 13 are unresolved point sources, 33 have extended emission that is consistent with star formation, and only 1 is identified as having jets and lobes\\footnote{2MASX\\,J18324117-3411274, which on visual inspection of the SUMSS image could also be confused with nearby point source companions.}. For the remaining 157 galaxies in our sample, which we classify as AGN dominated, 26 are unresolved point sources, 123 are identified as having jets and lobes, 3 have extended emission consistent with star formation\\footnote{NGC\\,2559, WKK\\,4748 and NGC\\,5903, all of which were not detected in either of the 60 and 100\\,$\\mu$m bands of \\emph{IRAS}.}, and 5 have unknown structure. Hence, there is a clear consistency between the classifications based on the radio and far-infrared luminosities and the radio and near-infrared morphologies. We summarize these classifications in Figs\\,\\ref{figure:luminosity_class} and \\ref{figure:redshift_class}, showing their distribution as a function of 1.4\\,GHz radio power and redshift.  For the sub-sample of star-forming galaxies, two factors significantly reduce the likelihood of detecting \\mbox{H\\,{\\sc i}} absorption in HIPASS -- the predominance of \\mbox{H\\,{\\sc i}} emission (which arises from the large reservoirs of gas required to form stars) and the distribution of the continuum flux density over the extended stellar disc. In the case of the former, the spatial distribution of the emission is typically unresolved by HIPASS and so acts to significantly mask any potential absorption of the background continuum at low redshifts. Furthermore, the continuum emission is extended over kpc scales, effectively reducing the covering factor $f$ and hence the likelihood of detecting absorption against a small fraction of the total flux density. If we assume that the sizes of cold and dense \\mbox{H\\,{\\sc i}} gas clouds are typically 100\\,pc \\citep[e.g][]{Braun:2012, Curran:2013b}, then the fraction of radio emission obscured by a single absorbing cloud will be $f \\sim 0.01$, which is equivalent to $S_{1.4}\\sim3$\\,mJy for the flux density limit of our sample, and so well below the noise level. While 33 of the star-forming galaxies are identified morphologically as having extended emission, some of the more compact point sources will have star formation concentrated within the sub-kpc nuclear region, effectively increasing the likelihood of absorption detection. This is certainly the case for the single detection of \\mbox{H\\,{\\sc i}} absorption we obtain in our sub-sample of star-forming galaxies, Arp\\,220, where a significant fraction of the nuclear radio emission is thought to be generated by starburst activity ($\\sim 240$\\,M$_{\\odot}$\\,yr$^{-1}$; \\citealt{Anantharamaiah:2000}).  A far smaller fraction of the AGN-dominated radio galaxies have \\mbox{H\\,{\\sc i}} emission detected in HIPASS compared with those that are star forming. This is in part due to their distribution towards higher redshifts, but also that many of these sources will be hosted by massive, neutral gas-poor, early-type galaxies (e.g. \\citealt{Bregman:1992}). However, while fewer have emission lines that could mask the detection of absorption at low redshifts, the majority have morphologies (jets and lobes) that are extended over scales greater than 45\\,arcsec (the typical spatial resolution of both NVSS and SUMSS), which at the median redshift of $cz = 6000$\\,km\\,s$^{-1}$ equates to physical scales greater than $\\sim 20$\\,kpc. The likelihood of absorption against these extended sources is low since most of the continuum emission will not be obscured by the discs or rings in which we expect the absorbing \\mbox{H\\,{\\sc i}} gas to be located. In the special case of Centaurus\\,A, the proximity of this radio galaxy to the Milky Way means that, while only a fraction of the total continuum emission is concentrated within the nucleus of the galaxy, we can still detect significant absorption of $\\Delta{S} \\approx 1$\\,Jy against the core. Furthermore, the \\mbox{H\\,{\\sc i}} emission and absorption are spatially resolved and so can be identified as separate components. If Centaurus\\,A were instead located at the sample median redshift of $cz = 6000$\\,km\\,s$^{-1}$, both the emission and absorption (which would decrease to less than 13\\,mJy) would no longer be detectable with HIPASS.  Our two remaining detections of absorption occur in AGN-dominated radio galaxies, 2MASX\\,J13084201-2422581 and NGC\\,5793, both with continuum radio emission that is compact with respect to the size of the stellar disc of the galaxy (see Fig.\\,\\ref{figure:image_overlays_cabb}). As we have already noted in Section\\,\\ref{section:individual_sources}, these galaxies are both classified as having edge-on disc morphologies (with inclinations of $i \\approx 75\\degr$) and Seyfert~2 AGN activity. It is towards these compact radio galaxies that we would expect to have the highest detection rate for absorption, where a significant fraction of the total continuum flux density will be absorbed by the chance alignment of foreground cold and dense \\mbox{H\\,{\\sc i}} gas. At higher redshifts, where the 21\\,cm emission line is not easily detectable in HIPASS, the compact and radio-loud nuclear starbursts will also contribute significantly to the detection rate, as was seen for Arp\\,220. If we consider just the point sources that do not have catalogued 21\\,cm line emission, then we obtain a detection rate for absorption of 6\\,per\\,cent (2/31), which is approaching the typical rates obtained by targeted searches of compact radio sources (see e.g. \\citealt{Allison:2012a} and references therein).  \\subsection{Comparison with the ALFALFA pilot survey}\\label{section:alfalfa_survey}  The ALFALFA survey on the Arecibo Telescope \\citep{Giovanelli:2005} is the only other existing large field-of-view survey for \\mbox{H\\,{\\sc i}} gas in the local Universe, which when completed will map $7000\\,\\deg^{2}$ of the sky in the redshift range $-2000 < cz < 19\\,000\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$. \\citet{Darling:2011} recently conducted a blind pilot survey of \\mbox{H\\,{\\sc i}} absorption in the volume bounded by $-650 < cz < 17\\,500\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ and $517\\,\\deg^{2}$ (1.3\\,per\\,cent of the celestial sphere and 7.4\\,per\\,cent of the full ALFALFA footprint). They found no intervening absorbers (which is consistent with the redshift search path and column density limits) and a single strong absorption line ($\\tau \\approx 0.64$) at $cz \\approx 10\\,800\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$, associated with the interacting luminous infrared galaxy UGC\\,6081 that had previously been detected by \\citet{Bothun:1983} and \\citet{Williams:1983}.  To compare their result with our HIPASS search, we again use the local radio luminosity function of \\cite{Mauch:2007} to estimate the expected number of galaxies above a flux density limit of $S_{1.4} = 42$\\,mJy (defined by a $5\\,\\sigma$ detection of absorption\\footnote{The per-channel noise in the smoothed ALFALFA spectra is 2.2\\,mJy.} with an optical depth of $\\tau > 0.3$), within the comoving volume bounded by 517\\,deg$^{2}$ of sky and the redshift range $0 < cz < 17\\,500\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ (approximately $8 \\times 10^{-4}$\\,Gpc$^{3}$). This yields approximately 29 galaxies, of which 19 are AGN dominated and 10 are star forming. The total detection rate based on this sample is therefore $\\sim 3$\\,per\\,cent (1/29), which is consistent with our results. We note that 10\\,per\\,cent of the redshift range was found to be unusable by \\cite{Darling:2011}, due to contamination from RFI and Galactic 21\\,cm emission, and so the expected detection rate is in fact slightly higher. UGC\\,6081 is not in the region of sky observed by \\emph{IRAS}, and so we cannot use the far-infrared versus radio luminosity relationship to classify this galaxy. The galaxy is in the process of a merger, exhibiting two radio nuclei that are separated by only 16\\,arcsec in the Faint Images of the Radio Sky at Twenty Centimeters survey (\\citealt{Becker:1995}), indicating that the radio emission may be arising from nuclear starburst activity in a similar mode to Arp\\,220. The radio power (assuming a total flux density of $S_{1.4} = 170$\\,mJy; \\citealt{White:1997}) at the redshift of the galaxy is $P_{1.4} \\approx 5 \\times 10^{23}$\\,W\\,Hz$^{-1}$, which is consistent with either a highly luminous starburst or AGN activity. UGC\\,6081 would likely be classified as either a compact AGN or nuclear starburst in our sample, which is consistent with the majority of our detections (excluding Centaurus\\,A as a special case).  \\subsection{Implications for future H\\,{\\sevensize\\bf I} absorption surveys}  \\begin{figure} \\centering \\includegraphics[width = 1.0\\columnwidth]{size_redshift.pdf} \\caption{Physical size versus redshift for the 15.5\\,arcmin HIPASS beam and the upper-limit 30\\,arcsec beam expected for a future \\mbox{H\\,{\\sc i}} absorption survey using ASKAP. The solid lines indicate the physical scale of the beam within the redshift range indicated by the vertical dashed lines. At the higher redshifts available to an \\mbox{H\\,{\\sc i}} absorption survey using ASKAP, a 30\\,arcsec beam will be sensitive to the physical scales resolved by HIPASS.}\\label{figure:size_redshift} \\end{figure}  We can use our results to estimate the number detections that might be achievable with the full ALFALFA survey. Considering a blind survey for associated absorption, conducted over 7000\\,deg$^{2}$ of the sky in the redshift range $0 < cz < 19\\,000$\\,km$^{-1}$ (equating to a comoving volume of 0.0136\\,Gpc$^{3}$), we use the local radio luminosity function of \\cite{Mauch:2007} to predict a total of 455 galaxies above a detection flux limit of 42\\,mJy (see Section\\,\\ref{section:alfalfa_survey} for an explanation of this limit). By simply applying our total detection rate from HIPASS, we predict approximately 10 detections of associated absorption, while applying the rate estimated for the pilot survey of \\cite{Darling:2011} yields approximately 16 (which is driven by the fractional increase in volume of the full survey). Since the ALFALFA survey probes higher redshifts than HIPASS, we expect that the detection rate amongst a radio flux density selected sample of galaxies will be higher, due to a decrease in the fraction of diffuse star-forming galaxies. We therefore predict a factor of 3--4 increase in the number of detections of associated absorption compared with what we have achieved in HIPASS.  Our results have positive implications for proposed all-sky absorption surveys on the precursor telescopes to the SKA, which will be able to probe the \\mbox{H\\,{\\sc i}} content of the Universe up to $z \\approx 1$. In Fig.\\,\\ref{figure:size_redshift}, we show that the physical scales to which HIPASS is sensitive in the local Universe are well matched at higher redshifts to the smaller beam sizes of the proposed First Large Absorption Survey in \\mbox{H\\,{\\sc i}} on the Australian Square Kilometre Array Pathfinder (ASKAP; \\citealt{Deboer:2009}). The HIPASS detections show that it will be possible to use ASKAP (with a spatial resolution of $d \\sim 100$\\,kpc at $z \\sim 0.5$) to detect strong absorption systems associated with compact AGN and nuclear starbursts at redshifts in the range $0.5 < z < 1.0$, probing an epoch of the Universe not yet explored by all-sky \\mbox{H\\,{\\sc i}} surveys.  Furthermore, our results indicate that such strong absorption systems detected in future all-sky surveys may well provide excellent targets for luminous H$_{2}$O megamaser detection at redshifts greater than 0.1, where so few have been discovered\\footnote{SDSS J0804+3607 at $z = 0.66$ \\citep{Barvainis:2005} and 4C\\,+05.19 at $z = 2.64$ \\citep{Impellizzeri:2008}, towards which multiple intervening and associated \\mbox{H\\,{\\sc i}} absorption lines have been detected \\citep{Moore:1999, Curran:2007b, Curran:2011c, Tanna:2013}.}. \\citet{Taylor:2002} showed that for the sample of known H$_{2}$O megamaser galaxies at that time, the detection rate of \\mbox{H\\,{\\sc i}} absorption was greater than 42\\,per\\,cent, higher than the typical rates achieved with targeted surveys of compact radio sources (see \\citealt{Allison:2012a} and references therein). Of the Seyfert~2 galaxies in which we have detected absorption, NGC\\,5793 has known H$_{2}$O megamaser emission from within a sub-pc disc around the AGN \\citep{Hagiwara:1997}, while 2MASX\\,J13084201-2422581, although not detected, has been a candidate for several large-scale megamaser surveys (e.g. \\citealt{Braatz:1996, Sato:2005, Kondratko:2006}). VLBI studies at high spatial resolution show that the optically thick \\mbox{H\\,{\\sc i}} absorption in these galaxies probably arises within an edge-on disc on 100\\,pc scales. If dense molecular gas also exists on sub-pc scales, it is likely to be distributed as a circumnuclear disc that is similarly orientated and so generate significant amplification of H$_{2}$O emission towards us, thereby providing favourable conditions for the detection of megamasers at cosmological redshifts. We also note that radio-loud nuclear starburst galaxies such as Arp\\,220, which is host to multiple regions of OH megamaser emission within its double nuclei \\citep{Lonsdale:1998, Rovilos:2003}, are likely to produce strong 1.6\\,GHz emission that can be detected in the wide frequency bands of these future \\mbox{H\\,{\\sc i}} surveys."
    },
    "1402/1402.5485.txt": {
        "abstract": "The increasing sensitivity of instruments at X-ray and TeV energies have revealed a large number of nebulae associated to bright pulsars. Despite this large data set, the observed pulsar wind nebulae (PWNe) do not show a uniform behavior and the main parameters driving features like luminosity, magnetization, and others are still not fully understood. To evaluate the possible existence of common evolutive trends and to link the characteristics of the nebula emission with those of the powering pulsar, we selected a sub-set of 10 TeV detections which are likely ascribed to young PWNe and model the spectral energy distribution with a time-dependent description of the nebulae's electron population. In 9 of these cases, a detailed PWNe model, using up-to-date multiwavelength information, is presented. The best-fit parameters of these nebula are discussed, together with the pulsar characteristics. We conclude that TeV PWNe are particle-dominated objects with large multiplicities, in general far from magnetic equipartition, and that relatively large photon field enhancements are required to explain the high level of Comptonized photons observed. We do not find significant correlations between the efficiencies of emission at different frequencies and the magnetization. The injection parameters do not appear to be particularly correlated with the pulsar properties either. We find that a normalized comparison of the SEDs (e.g., with the corresponding spin-down flux) at the same age significantly reduces the spectral distributions dispersion. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  During the last few years, the number of pulsar wind nebulae (PWNe) detected at TeV energies has increased from 1 (the Crab nebula, Weekes et al. 1989) to $\\sim$30. The latter number of detected PWNe, mostly contributed by the H.E.S.S. survey of the Galactic plane (see, e.g., Carrigan et al. 2013 for a recent status report), is similar to the number of characterized nebulae at other frequencies. The Cherenkov Telescope Array (Actis et al. 2011) will likely increase this number to several hundreds (de O\\~na Wilhelmi et al. 2013), probably providing an essentially complete account of TeV emitting PWNe in the Galaxy.  These recent PWNe discoveries provided a basic understanding of their phenomenology:  assuming that the PWNe is maintained solely by the pulsar rotational power,  the $\\gamma$-ray luminosity detected is believed to be the result of Comptonization of soft photon fields by relativistic electrons injected by the pulsar during its lifetime. This scenario can lead to TeV sources without counterparts (e.g., the first one was detected by Aharonian et al. 2002, Albert et al. 2008), when the synchrotron emission is reduced by the decay of the magnetic field. Also, it can lead to large mismatches in extension between $\\gamma$ and X-ray energies, when the magnetic field is low enough that electrons emiting keV photons  actually cool faster and are more energetic that electrons emitting in TeV (see de Jager \\& Djannati-Atai 2008 for a discussion). The explanation of these basic properties of the behavior of PWNe does not imply that we understand the population detected in detail.  \\subsection{Pulsars with low characteristic age}  A compilation of pulsars with known rotational parameters and characteristic age of $\\tau<10^4$ years is presented in { { Table 1, }} which is obtained from the updated ATNF catalog (Manchester et al. 2005) and includes the recently detected magnetar close to the Galactic Center (Mori et al. 2013, Rea et al. 2013). The value of the period $P$, period derivative $\\dot P$, distance $D$, characteristic age $\\tau$, dipolar field $B_d$, spin-down power $L_{sd}$, and $L_{sd} / D^2$ is listed. Their definitions are given below. These values are obtained directly from the catalog, neglecting some better estimations on the distances, such as those of e.g., G0.9+0.1 or pulsars at the LMC, in favor of uniformity when compiling the table. % According to their position in the sky, we added the label H, M or V (for H.E.S.S., MAGIC or Veritas respectively) to indicate the visibility from different Cherenkov telescopes. The names of the TeV putative PWNe (or at least co-located TeV sources even if the TeV source is likely not associated to the pulsar in some cases) are also included. The majority of these pulsars, located in the inner part of the Galaxy, were in the reach of the H.E.S.S. Galactic Plane Survey (GPS), which attains a roughly uniform sensitivity of 20 mCrab (Gast et al, 2012). Some of the pulsars in the northern sky have been observed by either MAGIC or Veritas, with comparable sensitivity.  To compare the pulsar sample in Table 1 we consider their characteristic ages. Even if this is not the pulsar real age, which is usually uncertain, it can be considered a good approximation when the pulsar braking index is $n \\sim 3$ and the initial pulsar spin-down period is much shorter than the current one. In order to give an idea of relative strength, the spin-down power of any pulsar is compared to that of the Crab extrapolated to the corresponding characteristic age.  The last three columns in Table 1  represent, respectively, the age of Crab (assuming no change in braking index) at which it would have the same characteristic age as the corresponding pulsar ($T^{Crab}_{\\tau}$), the Crab's spin-down power at that age ($ L_{sd}^{Crab}(T^{Crab}_{\\tau}) $), and the spin-down power of the pulsar in terms of percentage of $ L_{sd}^{Crab}(T^{Crab}_{\\tau} )$, which we refer to as CFP (or Crab fractional power). When looked in this way, the Crab pulsar is no longer special.    \\subsection{The influence of age in pulsars of similar spin-down}  Considering the characteristic ages provides the possibility of assessing the total power input into the nebula. Take as an example PSR J1617--5055 and J1513--5908, and assume for the sake of the argument that both generate TeV emission via a PWN. Both pulsars have essentially the same, and relatively high spin-down power, $1.7 \\times 10^{37}$ erg s$^{-1}$. However, one has likely been injecting this power for a much longer time, since the characteristic age of PSR J1617--5055  is a factor of 5 larger than that of PSR J1513--5908. The electrons that populate the nebulae will sustain energy losses and live, in most conditions, for more than 10$^4$ years. Thus, it is reasonable to suppose that there will be more high-energy electrons with which generate TeV radiation in the older pulsar than in the younger one. The differences between PSR J1617--5055 and J1513--5908 are reflected in the comparison with Crab at the moment when its characteristic age is correspondingly the same to the pulsar in question. PSR J1617--5055 is approximately three times as luminous than Crab will be at the same characteristic age. Instead PSR J1513--5908 spin-down corresponds to only a few percent of the one Crab will have at its characteristic age. Thus, even when both have the same spin-down we are speaking of very different nebulae.  \\begin{figure}[t] \\centering \\includegraphics[width=84mm]{taucomp.eps} \\caption{Comparing two PWNe models that differ only in age. Parameters of these models %, including $\\tau_0$, are as those used for the Crab nebula, and both have the same spin-down power.} \\label{taucomp} \\end{figure}  To exemplify further this point, consider two mock pulsars having the same spin-down evolution, magnetic fraction, injection spectrum, and photon background parameters than Crab (see below for precise definition of all these quantities) and both having also the same spin-down power, $1.7\\times10^{37}$ erg s$^{-1}$, but two different characteristic ages of 1500 and 8000 years, respectively. The modeled PWNe (details of the model itself are discussed below) when every parameter is the same but just the $\\tau$ and the corresponding real age vary % that is we assume the same tau_0 turn out to be different: For instance, the resulting magnetic field varies from 1 to 30 $\\mu$G. %and the size of the PWN changes from 2 to 40 pc. The SEDs shown in Fig. \\ref{taucomp} show that the spin-down power $L_{sd}$, or the parameter $L_{sd} / D^2$ (which is the same for the SEDs in the figure), unless of course when $L_{sd}$ is extremely low, cannot by themselves blindly define dectectability of PWNe, and further considerations about the PWNe age, injection, and environment have to be taken into account. This conclusion is emphasized when the photon background, the injection, and the magnetic fraction, among other key parameters, may vary from one pulsar to the next.    \\subsection{Recent models and differences}  Table 1 shows that most of the young pulsars we know of were indeed surveyed for TeV emission. This has motivated developing detailed radiative models to tackle the complexities in each of the PWNe. However, whereas some of these models are time-dependent, which is essential for a proper accounting of the nebula evolution and electron losses as per the discussion above, they are different to one another, and are constructed under different approximations and assumptions. Just considering the most recent literature, one can see that some models approximate the electron population computation to obtain an advective differential equation  (e.g., Tanaka \\& Takahara 2010; 2011), whereas others neglect the treatment of energy losses in full and instead replace it by the particle\u00d5s escape time (e.g. Zhang et al. 2008; Qiao, Fang, \\& Zhang 2009), and yet others do not impose any approximation at this level (e.g., Martin et al. 2012). Some models actually assume the particle population directly and neglect any time dependence in most of the magnitudes (e.g., Abdo et al. 2010). Some assume the injection is described by a broken power law (e.g., Bucciantini et al. 2011; Tanaka \\& Takahara 2010, 2011; Martin et al. 2012; Torres et al. 2013a,b), others consider that particle spectrum downstream of a relativistic shock can be fitted as a Maxwellian plus a power-law tail, despite the increased amount of unconstrained fitting parameters (e.g., Fang \\& Zhang 2010). Some impose conservation of the total energy injected by the pulsar summing up the energy fractions distributed in particles and magnetic field (e.g., Tanaka \\& Takahara 2011, Torres et al. 2013a,b); in others, this condition is relaxed  (e.g., Bucciantini et al. 2011). Some models have account of the dynamics beyond reverberation (e.g., Gelfand et al. 2009, Fang \\& Zhang 2010, Bucciantini et al. 2011), while most others do that with less precision. Some take into account self-synchrotron emission (e.g.,  Tanaka \\& Takahara 2011, Bucciantini et al. 2011, Torres et al. 2013b), others do not. Some consider bremsstrahlung (Martin et al. 2012), others do not, even when densities assumed are somewhat large (Li et al. 2010). Some models consider the magnetic field evolution by taking into account its work on the environment (e.g., Bucciantini et al. 2011; Torres et al. 2013a,b), others approximate it (e.g., Tanaka \\& Takahara 2010, 2011). The magnitude of spectral results introduced by different underlying assumptions has been quantified only in some cases (e.g., see the impact on approximating the electron computation in Martin et al. 2012). Having a clear conversion of results from one model to another, in order to generate a uniform theoretical setting where PWNe fittings can be compared, is simply impossible.        % In addition, apart from the obvious mismatches in the models per se, the nebulae that have been studied with each of them are scarce. Table \\ref{thmodels} gives some examples using a certainly incomplete span of the literature. Our interpretation of observations is based on uncommon modeling, undermining our conclusions.  \\subsection{This work}  The purpose of this work is to put at least a partial remedy to this situation, and provide a study of several young, TeV detected PWN. In order to do that we have improved our radiative model of PWNe (Martin et al. 2012) and applied it to observations. The model is one zone, leptonic, and time-dependent. It seeks a solution for the lepton distribution function considering the full time-energy-dependent diffusion-loss equation. The time-dependent lepton population is balanced by injection, energy losses and escape. We include losses by synchrotron, inverse-Compton (Klein Nishina inverse Compton with the cosmic-microwave background as well as with IR/optical photon fields), self-synchrotron Compton, and bremsstrahlung, devoid of any radiative approximations, and compute likewise the radiation produced by each process. We consider below in more detail the computation of the magnetic field evolution and its relation with the magnetization of the nebula. The main caveats of this model are that it contains only a free expansion dynamics (we come back to this below) and no geometry other than assuming spherical symmetry. These are clear over simplifications for some nebulae, where, for example, we know one size does not fit all frequencies. Still, it is a complete radiative model, and despite these caveats, it makes sense to use it for a more systematic study of the youngest nebulae.  Our sample is formed by 10 TeV detected, possibly Galactic PWNe, taken from Table 1 plus the recently detected CTA~1, which has a characteristic age slightly larger than $10^4$ years. In the Appendix of this work we comment on why we do not consider in our study the cases of HESS J1023--575, J1616--508, J1834--087/W41, and J1841--055 (in most cases, the information gathered on them imply  that the TeV emission is not univocally associated with a PWN) as well as Boomerang and HESSJ1640--465. We find that not all of the 10 cases studied are best interpreted with a PWN. In particular, we conclude that the case of HESS J1813--178 is most likely related to the SNR rather than to the PWN. The rest of this paper is organized as follows. The following Section briefly introduces the model used. Section 3 deals with each of the TeV detections in our sample, provide a PWN model when possible, and discussing the complexities of each case, surfacing caveats of our model when appropriate. Finally, Section 4 puts all our results in context, compare the modeled PWNe, and draws some conclusions from the overall population.       %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "The aim of this paper was to present numerical models of the TeV-detected, young PWNe along more than 20 decades of frequencies; using a radiatively complete, time-dependent numerical approach. For the first time, we have a coverage of many such PWNe analyzed under the same framework, adopting similar assumptions, which allows for a more meaningful parameter comparison. Despite the caveats of the model used, for instance, the simplicity of having adopted a uniform magnetic field, a free expansion phase, and of disregarding morphological features, we find that one-zone, leptonic-only generated radiation provides a reasonably good fit to the multifrequency data for PWNe detected at TeV. Here we summarize our findings.  \\begin{itemize}  \\item We favor a non-PWN origin for the radiation detected from HESS J1813--178. For the remaining 9 TeV sources studied, we find a plausible PWN origin of the mutiwavelength emission.  \\item For all the TeV sources  plausibly related with a PWN, only the Crab nebula is SSC dominated. All the remaining PWNe except for HESS J1356--645 and CTA 1 are IC-FIR dominated. The dominance of the FIR contribution to IC is always significant.  \\item The  FIR energy densities that we found is needed to fit the PWN high-energy emission are generally larger than what is obtained from GALPROP (usually by up to a factor of a few).  \\item The efficiencies of emission are $\\sim 10^{-6 \\div 7}$ in radio, $\\sim 10^{-2 \\div 3}$ in X-rays, and $\\sim 10^{-3 \\div 4}$ in gamma-rays, with only one outlier in the sample presenting very low X-ray fluxes (G292.2--0.5).  \\item The electron population can be described by a broken power law in all cases. The parameters of the injection cluster in relatively narrow ranges, especially, the break Lorenz factor, which is around $5 \\times 10^5$. The high energy spectral slope is found to be in the range 2.2 -- 2.8 (except for the steeper case of G292.2--0.5, which also present a higher energy break). The low energy part is instead much harder, with the low energy index in the range 1.0 -- 1.6.  \\item All PWNe have large multiplicities, in general in excess of $10^5$. The population of low-energy electrons is large by number, and generate a low medium energy per particle in the spectrum in all cases.  \\item All the nebulae except CTA~1 have low values of magnetization, of only a few percent. CTA~1 presents the largest magnetization of our sample, and reaches almost to equipartition. All the other PWNe are heavily particle dominated. This result is found to be stable against uncertainties.  \\item We do not find significant correlations between the efficiencies of emission at different frequencies and the magnetization, implying that the specific environment and the injection effects play a dominant role in determining, e.g.,  the gamma-ray luminosity.  \\item Comparing SEDs of the PWNe as observed today mixes pulsars of different spin-down power and age, and generates a variety of distributions. A normalized comparison of the SEDs (e.g., with the corresponding spin-down flux) at the same age significantly reduces the dispersion.  \\item We do not find clear correlations between the pulsar's characteristic ages and the radio and X-ray luminosities. The gamma-ray luminosity seems to be anti-correlated with the characteristic age. On the other hand, we do find correlations of the radio and X-ray (and at a slightly lower confidence also gamma-ray luminosities) with the spin-down,  and an anti-correlation of the ratios of IC to synchrotron luminosities with the spin-down.  \\item The injection parameters do not appear to be correlated with the pulsar properties, except for the maximum Lorentz factor and the magnetic field in the nebula which are correlated with the spin-down power (or the magnetic field at the light cylinder or the electric potential), but these cases can be ascribed to the model properties.  \\item We do not find a significant correlation of any PWN parameter with the surface magnetic field of the pulsars.  \\end{itemize}    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection*"
    },
    "1402/1402.0125_arXiv.txt": {
        "abstract": "We describe a methodology to classify periodic variable stars identified using photometric time-series measurements constructed from the Wide-field Infrared Survey Explorer (WISE) {\\it full-mission} single-exposure Source Databases. This will assist in the future construction of a WISE Variable Source Database that assigns variables to specific science classes as constrained by the WISE observing cadence with statistically meaningful classification probabilities. We have analyzed the WISE light curves of 8273 variable stars identified in previous optical variability surveys (MACHO, GCVS, and ASAS) and show that Fourier decomposition techniques can be extended into the mid-IR to assist with their classification. Combined with other periodic light-curve features, this sample is then used to train a machine-learned classifier based on the {\\it random forest} (RF) method. Consistent with previous classification studies of variable stars in general, the RF machine-learned classifier is superior to other methods in terms of accuracy, robustness against outliers, and relative immunity to features that carry little or redundant class information. For the three most common classes identified by WISE: Algols, RR Lyrae, and W Ursae Majoris type variables, we obtain classification efficiencies of 80.7\\%, 82.7\\%, and 84.5\\% respectively using cross-validation analyses, with 95\\% confidence intervals of approximately $\\pm$2\\%. These accuracies are achieved at purity (or reliability) levels of 88.5\\%, 96.2\\%, and 87.8\\% respectively, similar to that achieved in previous automated classification studies of periodic variable stars. ",
        "introduction": "The Wide-field Infrared Survey Explorer \\citep[WISE;][]{wright10} mapped the entire sky in four bands centered at wavelengths of 3.4, 4.6, 12, and 22 $\\mu$m (hereafter W1, W2, W3, and W4) from January 7, 2010 to February 1, 2011 that spanned both cryogenic and post-cryogenic phases. WISE conducted its survey from a sun-synchronous polar orbit using a 40 cm cryogenically cooled telescope equipped with four 1024 $\\times$ 1024 infrared (IR) array detectors that simultaneously imaged the same 47\\arcmin~$\\times$ 47\\arcmin~field-of-view in all bands. The WISE survey strategy alternated stepping the scan path forward and backward on subsequent orbits in an asymmetric pattern that approximately matched the $\\approx1\\arcdeg$ per day orbital precession rate. In this way, each point near the {\\it ecliptic plane} was observed on every other orbit, or every 191 minutes, yielding typically 12 independent exposures over one day. The number of exposures increases with ecliptic latitude, reaching over 6000 at the ecliptic poles.  WISE surveyed the sky approximately 1.5 times during its cryogenic phase that ended on September 29, 2010. Data continued to be collected for another four months to support the discovery of Near Earth Objects \\citep[the NEOWISE program;][]{Mainzer11}. During this post-cryogenic phase, 70\\% of the sky was scanned, with only the W1 and W2 detectors returning scientifically useful data. Overall, WISE covered the full sky slightly more than twice, with each sky-coverage separated by approximately six months. When all the mission data are combined, the median single-exposure depth-of-coverage on the ecliptic becomes $\\sim 24$ and the effective observation timespan (without the six month phase shift) $\\sim2$ days.  The WISE Single-Exposure Source Databases \\citep{cutri12} contain the photometry from each individual WISE exposure. These offer a unique opportunity to search for variable stars over the entire sky at wavelengths that are relatively immune to dust extinction. In particular, the most common pulsational variables (e.g., Cepheids, RR Lyrae, and Miras) have served as standard candles that enabled measurements of the size scale of the Milky Way and the Universe \\citep[e.g.,][]{tammann08}. Products from previous WISE data releases have already had a significant impact on the calibration of the RR Lyrae period-luminosity relation at mid-IR wavelengths \\citep{klein11,madore13,dambis14,klein14}. An all-sky census of pulsating variables offers an opportunity to improve our understanding of Milky Way tomography and the distribution of dark matter in the Galaxy through their association with relic streams from disrupted star clusters and dwarf galaxies \\citep{grillmair10}. Pulsational variables are also crucial for understanding stellar birth, structure, mass loss, and evolution \\citep{eyer08}. On the other hand, variability associated with eclipsing binaries (e.g., Algol, $\\beta$ Lyrae, and W Ursae Majoris types) provide laboratories for probing accretion, mass transfer, binary evolution, and exoplanets \\citep[for a review, see][]{percy07}.  The deluge of data from current and future time-domain surveys presents an enormous challenge for human-based vetting, classification, and follow-up. Fortunately, computers and efficient machine-learning (ML) algorithms are starting to revolutionize the taxonomic problem. A broad class of ML methods are generically referred to as ``supervised''. These methods entail defining a set of rules or models that best describe the relationship between the properties (features) of a data set and some known outcomes (e.g., classifications), and then using this ``trained'' mapping to predict the outcomes for new data. On the other hand, ``unsupervised'' methods attempt to blindly identify patterns and structures amongst the features of a data set, that is, with no pre-labelled classes or outcomes to learn from. For an overview of ML methods in general, see \\citet{hastie09}. For a review of ML applications in astronomy, see \\citet{ball10}. ML also provides a statistical framework to make probabilistic statements about the class(es) that a particular object with a set of observables or features could belong to. Given the same training model, these statements are also reproducible and deterministic, whereas a human-based classification approach is not.  Supervised ML techniques have gained considerable popularity in the automated classification of variable stars \\citep{wozniak04,eyer05,deboss07,mahabal08,blomme10,richards11}. In particular, \\citet{richards11} compared several classification frameworks for variable stars selected from multiple surveys and concluded that the {\\it random forest} ML classifier \\citep[a tree-based technique popularized by][]{breiman01} generally performed best. Due to their flexibility and robustness, {\\it random forests} have also attained popularity in the real-time discovery and classification of astronomical transients in general \\citep[e.g.,][]{bloom12,morgan12,brink13}. Guided by its success, we adopt the {\\it random forest} method as the primary ML classifier in this study.  The WISE All-Sky Release Source Catalog \\citep{cutri12} contains approximately 8 million sources where likely flux variables were flagged using a methodology similar to that described in \\cite{hoffman12}. We are in the process of constructing the WISE Variable Source Database (WVSDB) that builds upon this basic variability flagging in the Catalog. This database will contain confirmed variable sources with their light curves, derived properties such as periods and amplitudes, and where appropriate, the probabilities of belonging to specific known variability classes.  The first step towards the construction of the WVSDB is a framework that can automatically and probabilistically classify variables using features and diagnostics derived from the WISE single-exposure time-series measurements. This paper describes the methodology behind this framework, starting with the construction of a training (or ``truth'') sample leveraged on known variables classified in previous optical surveys and cross-matched to WISE to define robust mid-IR light curve features for classification. In particular, we show how Fourier decomposition techniques can be extended into the mid-IR to define the relevant features for discriminating between the various classes. This training sample is then used to fit and validate the popular {\\it random forest} machine-learning technique to assist with the future classification of WISE flux variables for the WVSDB.  This paper is organized as follows. In Section~\\ref{cscheme} we define the variability classes that WISE is most sensitive to. Section~\\ref{tsampf} describes the construction of the training sample and selection of the mid-IR light curve features. An analysis of these features for the various classes of interest is presented in Section~\\ref{fan}. The {\\it random forest} machine-learning method for automated classification is described and evaluated in Section~\\ref{ml}, where we also compare this method to other state-of-the-art machine-learning methods. Section~\\ref{cpred} gives an overview of the WVSDB, the classification plan, and how a feedback mechanism based on ``active-learning'' could be used to allow for selection biases in the input training sample. Results are summarized in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have described a framework to classify {\\it periodic} variable stars identified using metrics derived from photometric time-series data in the WISE single-exposure Source Databases. This framework will be used to construct an {\\it all-sky} database of variable stars in the mid-IR (the WVSDB), the first of its kind at these wavelengths. The reduced effects of dust-extinction will improve our understanding of Milky Way tomography, the distribution of dark matter, stellar structure, and evolution in a range of environments.  We identified several light-curve features to assist with the automated classification of WISE periodic variables, and found that Fourier decomposition techniques can be successfully extended into the mid-IR to define features for unambiguously classifying variable stars. Guided by previous automated classification studies of variable stars, we trained a machine-learned classifier based on the {\\it random forest} method to probabilistically classify objects in a seven-dimensional feature space. Random forests satisfy our needs in terms of flexibility, ability to capture complex patterns in the feature space, assessing feature importance, their relative immunity to outliers and redundant features, and for providing simple methodologies to support active-learning frameworks that can extend and refine training models to give more accurate classifications.  We constructed a training sample of 6620 periodic variables with classifications from previous optical variability surveys (MACHO, GCVS, and ASAS) and found  that the most common types that separate rather well and are reliably identified by WISE (given its sensitivity, observing cadences and time-spans) are Algols, RR Lyrae, and W Ursae Majoris type variables. This sample was used to construct an initial RF training model to assess classification performance in general and hence whether our method was suitable for constructing a WVSDB. From cross-validating a separate sample of 1653 pre-classified objects, our RF classifier achieves classification efficiencies of 80.7\\%, 82.7\\%, and 84.5\\% for Algols, RR Lyr, and W Uma types respectively, with 2-$\\sigma$ uncertainties of $\\sim \\pm$2\\%. These are achieved at purity (or reliability) levels of $\\gtrsim88\\%$ where the only source of ``impurity'' to each specific class is contamination from the other two contending classes. These estimates are similar to those of recent automated classification studies of periodic variable stars in the optical that also use RF classifiers.  Future work will consist of selecting good quality candidates for the WVSDB, the computation of light-curve features, further selection to retain periodic objects (above some statistical significance), then construction of the WVSDB. The three-class RF training model defined above will form the basis for initially predicting the classes with associated probabilities for the WVSDB. These probabilities will be thresholded to secure the ``winning'' classes. This first classification pass will inevitably leave us with a large fraction of unclassified objects. Our input training model has considerable room for expansion and improvement since during its construction, there were variable types that had to be removed since they were either too scarce or too close to an existing larger class to enable a reliable classification. Therefore, following the first classification pass, we will refine the training model using an active-learning approach (tied with contextual and/or follow-up information) where the improved statistics will enable us to better map and divide the feature space into more classes as well as sharpen the boundaries of existing ones. This will appropriately handle the ``known unknowns'', but WISE's all-sky coverage and sensitivity offers a unique opportunity to discover new and rare variable types, or new phenomena and sub-types in existing classes of pulsational variables and eclipsing binaries."
    },
    "1402/1402.4530_arXiv.txt": {
        "abstract": "Studies of Compton scattering by relativistic electrons in a strong magnetic field have been restricted to either incident photon angles $\\theta'$ aligned along the magnetic field $B$ or incident photon energies $\\omega'$ below the first pair creation threshold $\\omega'_{PC}$.  When these restrictions are relaxed there is a resonance in Compton scattering associated with pair creation (PC), that is analogous to but independent of known resonances associated with gyromagnetic absorption (GA).  As with the GA resonances, that may be labeled by the Landau quantum numbers of the relevant states, there is a sequence of PC resonances where the scattering cross section diverges.  In this paper, the lowest divergence is studied for incident photon energies satisfying ${\\omega'}^2\\sin^2\\theta'/(2eB)\\ll1$, assuming that the scattering electron is in its ground (Landau) state.  This lowest resonance affects only parallel-polarized photons. ",
        "introduction": "Compton scattering of relativistic electrons in a strong magnetic field has been investigated for over four decades (eg:\\cite{Ca71,H79,DH86,BAM86}).  Relatively recently \\citet{G00} extended the theory to magnetic fields in excess of the critical magnetic field, $B_{\\rm cr}=4.4\\times10^{13}$G.  Such ultrastrong magnetic fields are relevant to magnetars which include soft gamma repeaters and anomalous X-ray pulsars (\\cite{WT06,M08}). Gonthier et al. made several restrictive assumptions, including that the initial electron is in its ground (Landau) state, and that the initial photon is propagating parallel to the magnetic field. The first of these assumptions is well justified. An electron in a Landau state, $n>0$, relaxes in a very short time to its ground state, $n=0$, through gyromagnetic emission. The justification given for the second assumption can be summarized as follows: if photons are distributed over a wide range of angles in the laboratory frame, in which the electron is highly relativistic, then these photons are nearly all propagating at a small angle to the magnetic field in the electron's rest frame, by which is meant the frame in which its parallel momentum is zero, $p_z=0$. This small angle is assumed to be zero. However, the perpendicular momentum of the initial photon, which is denoted by $k'_\\perp=\\omega'\\sin\\theta'$ in natural units ($\\hbar,\\,c=1$), is an invariant, and so has the same value in both frames. Moreover, this invariant appears in the parameter $x'=k'^2_\\perp/2eB$ in the theory, and setting $x'=0$ is overly restrictive, in that it excludes a resonance that is associated with pair creation (PC). This resonance is in addition to a resonance associated with gyromagnetic absorption (GA), which is well known in the pulsar context as resonant Compton scattering (\\cite{DH86,D90,L96}). In this paper the assumption of parallel propagation for the initial photon is relaxed, and the scattering cross section including the PC resonance is considered.  The two Feynman diagrams for Compton scattering are presented in Figure~\\ref{fig:Feyndiag}, time increasing from right to left.  The solid lines represent particles (electrons or positrons) and the broken lines represent photons with wave 4-vectors ${k'}^\\mu=(\\omega',{\\bf k}')$ and $k^\\mu=(\\omega,{\\bf k})$.  Compton scattering is a second order process, depending on the fine structure constant squared, whereby an initial photon of energy $\\omega'$ is absorbed by an initial particle with energy $\\varepsilon$.  The particle jumps to an intermediate state of energy $\\varepsilon''$ and subsequently deexcites to a final state of energy $\\varepsilon'$ emitting a final photon of energy $\\omega$.  The energy of the particle is made up of its rest energy, its parallel momentum (parallel to the magnetic field upon which the $z$-axis is aligned) and its perpendicular momentum which is a product of the magnetic field strength and the particle's Landau quantum number.  For the initial particle, this is of the form \\begin{equation} \\varepsilon=\\sqrt{m^2\\,+\\,p_z^2\\,+2neB}. \\end{equation} The initial and final particles are taken to be electrons.  The intermediate particle can be either an electron or a positron and both these states are summed over as is their Landau quantum number $n''$.  The two types of resonance associated with Compton scattering appear as resonances in the first and the second Feynman amplitudes.  In the first Feynman diagram when the intermediate particle is an electron, the first type of resonance that occurs is the well known GA resonance.  Its divergence is usually handled by introducing a finite width, equal to the gyromagnetic decay width of the intermediate state, and adding it to the resonant denominator.  The PC resonance is associated with the second Feynman diagram, in which the intermediate particle is a positron.  Both of these resonances were identified as far back as 1979 by \\citet{H79}, who restricted his calculations to photon energies below the first PC threshold.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=8cm]{Weise_fig1.eps} \\caption{The two Feynman diagrams for Compton scattering in a magnetic field which have different labeling of the particle states to those in the absence of the field.  For example, for the initial particle one has $q=n,p_z,\\sigma$ where $n$ is the Landau quantum number, $p_z$ is the parallel momentum and $\\sigma$ is the spin.  Both the energy and parallel momentum are conserved.} \\label{fig:Feyndiag} \\end{center} \\end{figure}  At either a GA or a PC resonance the question arises as to whether the process should be regarded as a second-order process, corresponding to Compton scattering with a virtual intermediate state, or as a 2-step process, corresponding to two first-order processes with a real intermediate state. For the GA resonance, the first step is gyromagnetic absorption of the initial photon by the initial electron, with $n=0$, leaving an electron in an excited Landau level, $n''$. The second step is gyromagnetic emission of the final photon by this electron, with a transition back to the ground state $n'=0$. The lifetime of the intermediate state is determined by the (inverse of the) gyromagnetic decay probability. A conventional way of including this effect is through the natural line width, $\\Gamma$, of the intermediate state, such that the resonant denominator is replaced according to $\\omega+\\varepsilon'-\\varepsilon''\\to\\omega+\\varepsilon'-\\varepsilon''+i\\Gamma/2$. For the PC resonance, the two steps are pair creation of the final electron and the intermediate positron by the initial photon, and annihilation of the initial electron and the intermediate positron to produce the final photon. In this case, the lifetime of the intermediate state is determined by the (inverse of the) rate of creation of an electron-positron pair by the initial photon.  At magnetic field strengths far in excess of $B_{cr}$, the effects of photon dispersion and wave function renormalisation become important for parallel-polarized photons (eg:\\cite{HI13,SU10} and references therein).  This is clearly evident in the work by \\citet{CR09} who investigated the Compton effect in a strongly magnetized ($B=200B_{cr}$) medium consisting of both vacuum and plasma contributions.  The purpose of the present investigation is to compare Compton scattering for a range of incident photon angles under the same conditions as employed by \\citet{G00}, namely assuming the vacuum dispersion laws $k'^2=0$ and $k^2=0$ for the incident and final photons; Gonthier et al. only considered an initial photon angle $\\theta'$ of $180^\\circ$, defined as $0^\\circ$ in their paper.  Although only over a limited range of photon energies and for incident and final electrons in their ground state, the results of the current work cover the entire range of incident photon angles.  The magnetic field strengths considered are $0.1B_{cr},\\,B_{cr},\\,10B_{cr}$ and $100B_{cr}$, the same as those chosen in \\citet{G00}.  The notation used in this paper differs from that used by Gonthier et al. in that the incident photon's parameters are primed and the final photon's parameters are unprimed whereas in Gonthier et al.'s paper, they are the opposite.  The results found in the present paper show that, over the entire range of magnetic fields, the Compton scattering cross sections are greater in magnitude once the incident photon has a nonzero component of perpendicular momentum.  Further, at the two highest magnetic field strengths for which the first PC threshold falls within the range of the incident photon energies considered, the cross sections increase as the PC threshold is approached in a similar manner as when the first GA resonance is approached.  These results are relevant for the superstrong magnetic fields featured in magnetars, as discussed in \\citet{G00}, where the Compton scattering process remains an important energy loss process.  For example, \\citet{BH07}, \\citet{N08} and \\citet{Ba11} investigated resonant Compton scattering at the cyclotron frequencies as a mechanism behind the hard X-ray tails persistent in AXPs and magnetars. A related application is to the effect of Compton scattering on the shape of the radiation spectra of strongly magnetized neutron stars; for example, \\citet{Mu12} derived the kinetic equations for Compton scattering as a basis for radiation transport models.  As discussed in this paper, the resonance at the first PC threshold for oblique incident photon angles and ultrastrong magnetic fields can be at a lower energy than the GA resonance, and give a previously unrecognized form of resonant Compton scattering in these applications.  The paper is set out as follows.  The relevant equations and approximations used in the evaluation of the Compton scattering cross section are set out in Sect.~2.  A detailed derivation of the second order $S$-matrix for Compton scattering from first principles is presented in this section.  The criterion for the intermediate state to be part of a 2-step process, rather that a virtual state in a scattering event, is discussed in Sect.~3.  In Sect.~4, specific details are given of how the PC resonance is handled.  The results are presented in Sect.~5 over a restricted region of incident photon energies below the first GA resonance.  Some discussion of the work by Gonthier et al. is given in this section.  In Sect.~6, a general summary of the important features of the paper is presented.  Natural units are used throughout. ",
        "conclusions": "The pair creation divergence present in Compton scattering was recognized by Herold in 1979 and it has been discussed by a number of authors (eg: \\cite{NK93},\\cite{GHS95},\\cite{K96}).  The two most recent papers that attempted to tackle the divergence formally were the one by \\citet{GHS95} and the other by \\citet{K96}.  In the former paper, the authors abandoned energy conservation entirely and assigned a Weisskopf-Wigner decay width to the initial and final electrons and photons as well as the intermediate particle.  This yielded a scattering matrix element with a time dependence representing the time difference between the formation of the initial states and the measurement of the final states.  In the latter paper, Kachelriess introduced Licht fields and spectral functions.  The main point made in the present paper is that the PC resonance is analogous to the more familiar GA resonances in ``resonant Compton scattering.'' At the resonance, the scattering formally separates into a two-stage process, gyromagnetic absorption and emission at a GA resonance, and pair creation and annihilation at a PC resonance. The usual procedure of including the broadening of the resonance due to the finite lifetime of the intermediate state applies to both resonances, with the relevant lifetime at a PC resonance involving that due to pair creation.  In summary, due to the restriction imposed on $x',x$ and $n,\\,n'$, only $\\omega'$ values below the GA resonance are considered in this study.  Nevertheless, some important results are obtained when $\\theta'$ is not aligned along the magnetic field.  In the energy regions where the PC resonance does not occur, the Compton scattering cross sections are larger by one to two orders of magnitude at the lower frequencies, irrespective of the strength of the magnetic field.  There is additional strength at the two highest magnetic fields at frequencies in the vicinity of the PC resonance.  Unlike the GA resonance which is approximately symmetric and of a finite width affecting all 4 possible polarization mode combinations, the lowest PC resonance has a strong broad tail and affects only the $\\parallel'\\,\\parallel$ mode.  (Higher PC resonances affect the other modes.)  The PC divergence however is only apparent when the incident photon is not propagating parallel to the magnetic field.  The photon's angle of incidence need only be slightly oblique for the resonance associated with pair creation to dominate the Compton scattering cross section.  Under these conditions, Compton scattering becomes an important process for the energy conversion of parallel-polarized photons.  The PC resonance in Compton scattering needs to be taken into account when considering the propagation of high energy ($\\gg$ few MeV) photons in superstrong magnetic fields ($B\\gg B_{cr}$) in magnetars."
    },
    "1402/1402.6159_arXiv.txt": {
        "abstract": "We present INTEGRAL spectral analysis in the orbital/superorbital phase space of \\lsi\\/. A hard X-ray spectrum with no cutoff is observed at all orbital/superorbital phases. The hard X-ray index is found to be uncorrelated with the radio index (non-simultaneously) measured at the same orbital and superorbital phases. In particular, the absence of an X-ray spectrum softening during the periods of negative radio index does not favor a simple interpretation of the radio index variations in terms of changes of state in a microquasar. We uncover hints for the superorbital variability in the hard X-ray flux,  in phase with the superorbital modulation in soft X-rays. An orbital phase drift of radio peak flux and index along the superorbital period is observed in the radio data. We explore its influence on a previously reported double peak structure of radio orbital lightcurve, posing it as a plausible explanation. ",
        "introduction": "\\lsi\\/ is a high mass X-ray binary hosting a B0 Ve star with an equatorial outflowing disk. The companion is orbited by a compact object with an orbital period of 26.496 days (e.g., Casares et al. 2005). \\lsi\\/ is detected up to GeV and TeV gamma-rays (Albert et al. 2006; Abdo et al. 2009; Acciari et al. 2009; Hadasch et al. 2012), and the non-thermal spectral energy distribution is dominated by high energy photons. Two short ($<$ 0.1 s) and highly luminous ($ > 10^{37}$ erg s$^{-1}$) thermal flares were detected recently from the direction of \\lsi\\/ (Barthelmy et al. 2008; Burrows et al. 2012), giving support to the hypothesis that the compact object in \\lsi\\ is a (at least internally) highly magnetized neutron star (Torres et al. 2012; Papitto et al. 2012).  Besides the orbital period of 26.496 days, and other short-timescale variability (see e.g., Paredes et al. 2007), \\lsi\\ is known to have a long-term (1667 days) superorbital modulation. The latter was first discovered in radio (Paredes 1987; Gregory 1989; Gregory 1999; Gregory 2002), and then observed in H${\\alpha}$ emission line (Zamanov et al. 1999), X-rays (Li et al. 2012; Chernyakova et al. 2012), GeV (Ackermann et al. 2013), and hinted at TeV (Li et al. 2012; Torres et al. 2012). The possible origin of this 1667 days superorbital modulation could be the precession of the Be disk (Lipunov \\& Nazin 1994), the beat frequency between the orbital and precessional rates in a microquasar scenario (Massi \\& Jaron 2013), or the cyclic variability in the Be star envelope. The last interpretation seems the most natural, since the H${\\alpha}$ emission line varies on the same period and it is well known for the Be stars that the size of the circumstellar disk grows as the equivalent width of H${\\alpha}$  increases (e.g., Hanushik, Kozok \\& Kaizer, 1988). In the maximum of the equivalent width of H${\\alpha}$, the X-ray as well as  $\\gamma$-ray emission are enhanced (Li et al. 2012; Ackermann et al. 2013). The orbital phase of the X-ray peaks from \\lsi\\ varies from 0.35 to 0.75 along the superorbital period, consistently leading the radio peaks by 0.2 orbital phase (Chernyakova et al. 2012). The different origin regions of X-ray and radio emission could lead to the observed phase lag. The observed multiwavelength variabilities in superorbital timescale were discussed in the scenario of a neutron star switching between ejector and propeller states by Torres et al. (2012) and Papitto et al. (2012).      Black hole composed microquasars manifest themselves in five classical spectral X-ray states  (van der Klis 1994; Tanaka \\& Lewin 1995; Tanaka \\& Shibazaki 1996; Esin et al. 1997): the high/soft (HS) state, the low/hard (LH) state, the quiescent state, the very high (VH) state, and the intermediate state. These spectral states were regrouped by McClintock \\& Remillard (2006) as thermal--dominant (TD), hard X-ray, quiescent, steep powerlaw (SPL), and intermediate states. From an observational perspective, an optically thick steady jet with flat (radio spectral index $\\alpha \\sim 0$) or inverted ($\\alpha \\geqslant 0$) radio component is usually seen when in the hard X-ray state (a power law spectrum with index 1.5 $< \\Gamma <$ 2.1 in X-rays, McClintock \\& Remillard 2006), and an optically thin transient jet ($\\alpha < 0$) is associated with the overall hard to soft state transition (Fender et al. 2004; Fender et al. 2009).  In the hypothesis of \\lsi\\ being a microquasar, the explanation for the short thermal flares that have been detected from the region of \\lsi\\ must be the spatial superposition of a magnetar with a gamma-ray binary, however unlikely. In this hypothesis, Zimmermann \\& Massi (2012) expected a steady jet ($\\alpha \\geqslant 0$) in the hard X-ray state of the source, thus characterized by a power law with index 1.5 $< \\Gamma <$ 1.8 and a cutoff at high energies. A transient jet (with $\\alpha < 0$) would on the contrary be expected in the SPL state of \\lsi, showing an X-ray slope of $ \\Gamma \\geq $ 2.4. These seem very simplified assumptions, given that such an SPL state is poorly characterized. McClintock \\& Remillard (2006) note that many BH binaries could be radio quiet during this phase, and moreover, such state seems often described by high X-ray luminosity ($> 0.2$ times the Eddington level) and the detection of a thermal component or QPOs; all of which are clearly not the case of \\lsi\\ (Li et al. 2011).     Although GBI radio data showed spectral index transitions (Massi \\& Kaufman Bernad\\'{o}, 2009), the photon index of INTEGRAL observations was found to be compatible with a low hard state (Chernyakova et al. 2006; Zhang et al. 2010) at all times. Only in one occasion the photon index was $\\Gamma=3.6^{+1.6}_{-1.1}$ (Chernyakova et al. 2006). However, the large error bars (due to little exposure) introduced a large uncertainty. Zimmermann \\& Massi (2012) demonstrated that averaging the INTEGRAL data over too large orbital and superorbital phase ranges can result in a dominant low hard state and smear out the SPL state. This effect, they posed, is what explains the hard spectra all along the orbit in the \\lsi\\ system, and which led to the inconsistence between data and their predictions. Here, using INTEGRAL observations that double the size of our previous analysis (Zhang et al. 2010) and two years more data than in Chernyakova et al. (2012) we can separate in orbital and superorbital phases and test these ideas. ",
        "conclusions": "We have tested one of the methods proposed to identify the microquasar nature of the \\lsi\\ system: Namely, that it is in SPL state when the radio index $\\alpha < 0$ and in low hard state when $\\alpha > 0$, and that using large orbital and superorbital phase ranges mixes these states smearing out the state transitions signatures (Zimmermann \\& Massi 2012). The latter has been claimed as the reason of the non-detection of significant spectral variations in the INTEGRAL data analyzed so far. Here, with a significantly increased dataset,  we have introduced the orbital / superorbital phase space and extracted the X-ray spectrum for the corresponding specific cuts, avoiding averaging data over a too large superorbital range. \\lsi\\ shows a hard spectrum (single power law with no cutoff, $ \\Gamma $=1.49$^{+0.21}_{-0.19}$, see Table 1) when $\\alpha < 0$, which is inconsistent with the SPL state interpretation of Zimmermann \\& Massi (2012) stated above.   As commented in the introduction, such a clear cut distinction between states based on the radio index seems unlikely. The SPL state itself is not an ``intermediate state\"  that lies between soft and hard states, though transitions pass through SPL frequently (McClintock \\& Remillard 2006; Zhang et al. 2013). It may also be unrealistic to think that optically thin radio emission is confined to the SPL state only. If a hard to soft state transition does not pass through an SPL state (e.g., the 2011 outburst of H 1743-322 during which the photon index is always below 2.4, Zhou et al. 2012) or optically thin radio flares appear in other intermediate states but not in SPL state, the X-ray photon index will not go above 2.4 during optically thin radio emission. Thus, the X-ray spectrum during $\\alpha < 0$ might be a mixture of a SPL state or transitions to SPL states (e.g. see the cases of XTE J1650-500, Corbel et al. 2004; GX 339-4, Gallo et al. 2004; XTE J1859+226, Brocksopp et al. 2002, Corbel et al. 2004), and intermediate states when optically thin radio emissions locate out of SPL states. Nevertheless, the fact that most of the observation time in INTEGRAL indeed has a concurrent $\\alpha < 0$ (at the same orbital and superorbital phase) indicates that the most likely alternative is simply that there is no state transition at all, something also supported by the low X-ray flux at all times.   \\begin{figure*}[t!] \\centering \\includegraphics[angle=0, scale=0.8] {two_peak.eps} \\caption{Left {and Middle} Panel: lightcurves of 2.25 GHz (top panel), 8.3 GHz (middle panel) and radio index (bottom panel) during superorbital phase 0.0--0.1. {Superorbital phase is shown as the upper axis on top of the panels}. Right Panel: folded orbital lightcurve of 2.25 GHz (top panel), 8.3 GHz (middle panel) and radio index (bottom panel) during superorbital phase 0.0--0.1. The two peaks reported in Massi \\& Kaufman Bernad\\'{o} 2009 are shown with dashed black lines. Different orbit (orbital phase 0.0--1.0) are shown with different color.} \\label{mos-1} \\end{figure*}     We have also shown hints for the superorbital variability of \\lsi\\ in hard X-rays, and that the appearance of a double peak structure in the radio lightcurve of superobitally folded data can be accommodated by a drift of a single peak in an orbital basis."
    },
    "1402/1402.2498_arXiv.txt": {
        "abstract": "Recently, both theoretical arguments and observational evidence suggested that a small fraction of fast radio bursts (FRBs) could be associated with gamma-ray bursts (GRBs). If such FRB/GRB association systems are commonly detected in the future, the combination of dispersion measure (DM) derived from FRBs and redshifts derived from GRBs makes these systems a plausible tool to conduct cosmography. We quantify uncertainties in deriving the redshift-dependent ${\\rm DM_{\\rm IGM}}$ as a function of $z$, and test how well dark energy models can be constrained with Monte Carlo simulations. We show that with potentially several 10s of FRB/GRB systems detected in a decade or so, one may reach reasonable constraints on $w$CDM models. When combined with SN Ia data, unprecedented constraints on dark energy equation of state may be achieved, thanks to the prospects of detecting FRB/GRB systems at relatively high redshifts. The ratio between the mean value $\\left< {\\rm {DM}_{\\rm IGM}} (z)\\right>$ and luminosity distance ($D_{\\rm L} (z)$) is insensitive to dark energy models. This gives the prospects of applying SN Ia data to calibrate $\\left< {\\rm {DM}_{\\rm IGM}} (z)\\right>$ using a relatively small sample of FRB/GRB systems, allowing a reliable constraint on the baryon inhomogeneity distribution as a function of redshift. The methodology developed in this paper can also be applied, if the FRB redshifts can be measured by other means. Some caveats of putting this method into practice are also discussed. ",
        "introduction": "The nature of late time cosmic acceleration is a deep mystery in cosmology and fundamental physics, which could be explained by introducing an exotic form of energy content with negative pressure, dubbed dark energy. Cosmological parameters have been measured via various standard candles or rulers, such as Type Ia supernovae  (SN Ia) \\citep{riess98}, baryon acoustic oscillations (BAO) \\citep{anderson12,beutler11}, as well as small scale anisotropies of the cosmic microwave background (CMB) radiation \\citep{hinshaw13,planck13}. Being bright beacons from deep universe, gamma-ray bursts (GRBs) have been considered as a potential complementary probe to conduct cosmography. Many authors have made use of GRB luminosity indicators as standard candles \\citep[e.g.][]{dai04,ghirlanda04,liangzhang05,schaefer07}. However, unlike the SN Ia candle, GRB correlations lack physical motivation and usually have relatively large scatter, so that their role as standard candles is debated. Nonetheless, GRBs can serve as a complementary tool to probe the relatively high-$z$ universe, such as the star formation history, the metal enrichment history, and the properties of intergalactic medium \\citep[e.g.][]{barkana04,mcQuinn08,virgili11,wang12}.  Recently, \\cite{thornton13} reported the discovery of a new type of cosmological transients, dubbed Fast Radio Bursts (FRBs). These objects have anomalously high dispersion measure (DM) values corresponding to a cosmological redshift between 0.5 and 1 \\citep{lorimer07, thornton13}\\footnote{Some recent works also invoke a galactic origin for some of the FRBs \\citep{loeb14,bannister14,kulkarni14}.}. If the redshifts of these events can be measured, the combination of $z$ and DM information would be invaluable to conduct cosmography. \\cite{zhang14} suggested that a small fraction of FRBs could be physically connected to some GRBs, whose central engine is a supra-massive millisecond magnetar which collapses to a black hole at $10^2 - 10^4$ seconds after the burst\\footnote{Most FRBs would be produced by supra-massive NSs collapsing into a black hole after a much longer delay, e.g. thousands to millions of years \\citep{falcke14}.}. Two possible such associations might have been observed by \\cite{bannister12}, and the fraction of GRBs that might host a magnetar central engine, and hence, a possible FRB, could be up to $\\sim 60\\%$ for long GRBs \\citep{lu14} and probably an even higher fraction for short GRBs \\citep{rowlinson13}. Using the dispersion measure (DM) values of the two possible FRB/GRB association candidates \\citep{bannister12}, \\cite{deng14} derived the upper limits of the baryon mass density along the line-of-sight of the two GRBs, which are consistent with the values derived by other methods. This lends further support to FRB/GRB associations. The FRB/GRB association systems, if commonly detected in the future, would be an ideal tool to constrain cosmological parameters and properties of dark energy at redshifts not attainable by SN Ia. We term this prospect as ``FRB/GRB cosmography''.  In this paper, we study the prospects of conducting FRB/GRB cosmography in detail. An independent work was recently carried out by \\cite{zhou14}, who discussed using FRBs to constrain dark energy properties assuming that the redshifts of FRBs can be measured. ",
        "conclusions": "FRB/GRB systems, if confirmed to be commonly exist, have great potential to infer cosmological parameters, especially to constrain the equation of state of dark energy. We have shown that with a moderate sample size of several 10s, one may reach a constraint on $w$ comparable to a large SN Ia sample. Combining SN Ia data and FRB/GRB systems, one may achieve unprecedented accuracy in constraining $w$. The insensitivity of the $D_{\\rm L}(z) / \\left<{\\rm DM}_{\\rm IGM}\\right> (z)$ ratio offers the advantage of using a relatively small sample to cabibrate $\\left<{\\rm DM}_{\\rm IGM}\\right> (z)$ and to diagnose the local IGM inhomogeneity as well as its redshift distribution.  Our method is also applicable if the redshifts of FRBs can be measured independently with other methods. \\cite{zhou14} explored this possibility and reached the similar conclusion that FRBs can constrain dark energy equation of state and IGM inhomogeneity. FRB/GRB associations provide a practical method to measure redshifts of FRBs, and the sample size adopted in this paper is based on observational and theoretical insights of GRBs as well as the results of preliminary searches of FRB/GRB associations \\citep{bannister12}.  In the end, we want to point out some caveats of our method:  1) If most of the FRBs (or even all of them) are eventually proved to be of a galactic origin, or the redshifts of FRBs could never be measured, FRBs are no longer relevant for cosmography.  2) As shown in our simulations, the constraint results for cosmological parameters are most sensitive to the intrinsic IGM inhomogeneity $\\sigma_{\\rm IGM}$. Even though the calibration of $\\left<{\\rm DM}_{\\rm IGM}\\right>$ could be achieved (see section 4), if $\\sigma_{\\rm IGM}$ is very large, its uncertainty could still reduce the constraint accuracy dramatically. A (much) larger sample size than simulated here is needed to make this method competitive.  3) The GRB hosts likely evolve with redshift \\citep[e.g.][]{perley13}. This suggests that both $\\rm{DM_{host}}$ and $\\rm{\\sigma_{host}}$ could be redshift dependent. In our simulations, we adopt a constant free parameter for $\\rm{\\sigma_{host}}$, and its contribution to $\\sigma_{\\rm DM_{IGM}}$ is reduced at higher redshifts due to $(1+z)$ factor (Eq.(\\ref{eq:sigma}). If it turns out that $\\rm{DM_{host}}$ and $\\rm{\\sigma_{host}}$ would increase with redshift and $\\rm{\\sigma_{host}}$ could become comparable or even larger than $\\sigma_{\\rm IGM}(z)$, the accuracy of our results could be reduced dramatically. To keep the method competitive, either a larger sample is required, or more observational information about the GRB host galaxy population is needed.  4) If the dispersions of various parameters is intrinsically non-Gaussian, a larger sample than simulated may be needed to quantify these dispersions to achieve an accurate constraint on dark energy models.   \\vskip 0.2in  We thank Fa-Yin Wang for a helpful discussion. This work is partially supported by National Basic Research Program (``973\" Program) of China under Grant No. 2014CB845800, the NSFC (11273005) and SRFDP (20120001110064).  HG acknowledges a Fellowship from China Scholarship Program for support."
    },
    "1402/1402.2983_arXiv.txt": {
        "abstract": "We present a multiwavelength analysis of the nebula around the candidate luminous blue variable G79.29+0.46. The study is based on our radio observations performed at the Expanded Very Large Array and at the Green Bank Telescope and on archival infrared datasets, including recent images obtained by the \\emph{Herschel} Space Observatory.  We confirm that the radio central object is characterized by a stellar wind and derive a current mass-loss rate of about $1.4\\times10^{-6}\\, \\rm M_{\\odot} yr^{-1}$. We find the presence of a dusty compact envelope close to the star, with a temperature between $\\sim$ 40 and 1200 K.  We estimate for the outer ejecta an ionised gas mass of 1.51 $\\rm M_{\\odot}$ and a warm (60--85 K) dust mass of 0.02 $\\rm M_{\\odot}$. Diagnostics of the far-infrared spectra indicate the presence of a photo-dissociation region around the ionised gas.  Finally, we model the nebula with the photo-ionization code CLOUDY, using as input parameters those estimated from our analysis. We find for the central star a luminosity of 10$^{5.4} \\rm L_{\\odot}$ and an effective temperature of 20.4 kK. ",
        "introduction": "Evolved luminous stars with an initial mass more than 20 $\\rm M_{\\odot}$ are believed to enter an instability phase, historically known as luminous blue variable (LBV).  This phase is characterized by intense mass-loss rates ($\\sim10^{-5}-10^{-3}\\, \\rm M_{\\odot} yr^{-1}$), associated to visual/near-IR spectroscopic and photometric variabilities \\citep{1994HD}, before the star evolves as Wolf-Rayet (WR), once it has reduced its mass \\citep{1976Conti}. However, the occurrence of this phase during the massive stars evolution and the mass-loss mechanism are still not well understood, mostly because of the rarity of LBVs in our Galaxy. For the moment, the objects that satisfy the variability criteria coupled with high mass-loss rates are only 12 in the Milky Way,   but recently the number of candidate LBVs (cLBVs) has been increased to 55 \\citep{1994HD, vanGenderen2001, Clark2005, 2010Gvaramadze, 2011Wachter, 2012Naze}, based on the discovery of ring nebulae surrounding these luminous stars. In fact, as a result of the mass-loss, such evolved massive stars can be surrounded by circumstellar nebulae, rich of processed material \\citep{2001Lamers}. LBV nebulae (LBVNe) have been observed in about 20 objects among LBVs and cLBVs in our Galaxy and they have shown to be dusty, hence emitting in a wide spectral domain. Such nebulae are indeed the fingerprints of the mass-loss phenomenon, therefore multiwavelength observations provide unique insights in the LBV mass-loss history.  The importance of studying LBVNe is increased by the possibility that they can skip the WR stage and experience the explosive supernova (SN) event \\citep[e.g. ][]{2006KotakVink,2008Vinkbis, 2013vanDyk}. If this is true, the LBV mass-loss phenomenon is of great interest also for the study of type IIn SNe, whose evolution is influenced by the circumstellar material originated through mass-loss of the progenitor star \\citep{2003Chugai, 2007SmithN, 2013Ofek}. The mass-loss archaeology appears to be, therefore, fundamental for different aspects of the massive stars evolution.  In this work we propose a new study of the ejecta around G79.29+0.46, a galactic candidate LBV located in the Cygnus-X \\citep[$D\\sim1.7$ kpc,][]{2006Schn}, one of the richest star forming regions in our Galaxy. Thanks to its location and also to its beautiful aspect and intriguing nature, this star and its associated nebula have been widely observed in spectral domains ranging from the optical to the radio. These observations have allowed to investigate different components of the ejecta associated to G79.29+0.46, which showed to have lost mass with different events \\citep{2010Kraemer, 2010Esteban}. Diverse works focused on this source can be mentioned \\citep[e.g.][]{1994Higgs, 1996Waters, 2000Voors, 2008Rizzo, 2011UmanaG79}, but a comprehensive study of the star and its associated nebula has never been performed as yet. Moreover, the physical parameters determined in these works are often based on different assumptions, making difficult their comparison and interpretation. For instance, \\citet{1994Higgs} interpreted the ring-like shell around the star as interstellar material swept-up  by the stellar wind and estimated for this shell a mass of about 14-15 $\\rm\\,M_{\\odot}$. Later, the nebula has been recognized as ejecta associated to the star \\citep{1996Waters}, but a reappraisal of its properties is still missing. The nebular mass inferred by \\citet{1994Higgs} and \\citet{1996Waters} suggests that the nebula associated to G79.29+0.46 is amongst the most massive known \\citep[cf. Table 9 in][]{2003Clark}. If correct, the mass would rival that of the Homunculus around $\\eta$ Carina \\citep[e.g.][]{1999Morris,2003Smithhomu}, but the central star of G79.29+0.46 is likely over an order of magnitude fainter \\citep[about $10^{5.5}\\rm\\,L_{\\odot}$, while $\\eta$ Car is $10^{6.7}\\rm\\,L_{\\odot}$,][]{1994Higgs,2000Voors, 2001Hillier} and consequently far less massive. Therefore, the nebular mass would represent a higher proportion of its initial mass.  In this paper we propose a new analysis of the multiple components of the nebula around G79.29+0.46, by means of recent observations which we performed in the radio domain, new far-IR images provided by the \\emph{Herschel} Space Observatory and other archival far- and mid-IR spectroscopic and photometric datasets collected in the last 20 years by different telescopes. The main assumption in this work is that the photometric and spectroscopic properties of this star has not changed significantly in the considered time-lapse, hypothesis reinforced by the fact that G79.29+0.46 is still a candidate LBV, since no significant visual or near-IR variability has been observed so far \\citep{2000Voors, 2008Vink}.  This work started with our observations of G79.29+0.46 performed at the Expanded Very Large Array (EVLA). Preliminary results of these observations were  presented in \\citet{2011UmanaG79} (Paper 1 hereafter), and the highest resolution radio image was compared with the \\emph{Spitzer} infrared images, revealing us that the multiple shells around the star have signature of different populations of dust and that the radio emission is ionization bounded. In this paper we complement the preliminary work in Paper 1 with new radio observations performed at the Green Bank Telescope (GBT) to overcome the typical problem of interferometers, which is the ``zero baseline missing\", resulting in the impossibility to determine the total flux of extended sources. Therefore we propose a new analysis of the radio emission aimed at determining important physical parameters related to the ionized gas and to the central star.  We also perform an analysis of the IR datasets retrieved from the data archives, in order to determine  IR lines and physical properties related to the dust component.  The paper is organized as in the following: in Section \\ref{par:obs} we describe our observations and the archival  datasets considered in this study. Then we present our analysis of the radio datasets (Section \\ref{par:maps}), focused on the ionized gas distribution and emission, while in Section \\ref{par:infrared} we study the dust emission and properties, probed by the IR images. Section \\ref{sec:irspectra} concerns the spectroscopic mid- and far-IR analysis of the nebula and in Section \\ref{sec:sed} we present a simple model obtained with the photo-ionization code Cloudy. We conclude with a brief description of the star mass-loss history in Section \\ref{sec:discus}. ",
        "conclusions": "\\label{sec:discus} We have proposed a multiwavelength analysis of the nebula around the LBV candidate G79.29+0.46 by considering data in a wide spectral range (from the radio to the mid-IR). The aim of this work was to explain all the observed properties under the same hypothesis and to provide a comprehensive picture of the nebular properties.  Summarizing:\\\\ 1) radio observations and the model obtained with the code Cloudy have allowed as to constrain the physical parameters of the star, which has a luminosity of $\\sim10^{5.4}\\,\\rm L_{\\odot}$ and an effective temperature of $20.4\\times10^3\\,\\rm K$;\\\\ 2) the star is currently loosing mass through stellar winds (as provided by analysis of the central object spectral index) with a rate of $1.4\\times10^{-6}(\\frac{v_{\\infty}}{110\\,\\rm km\\,s^{-1}})\\ (\\ \\frac{D}{1.7\\,\\rm kpc})\\ ^{3/2} f^{1/2}\\,\\rm M_{\\odot}\\,yr^{-1}$;\\\\ 3) analysis of the IR emission of the central object reveals that the star is surrounded by a close dusty envelope with temperature in the range from $\\sim$ 40 to $\\sim$ 1200 K; \\\\ 4) at a distance of 0.66 pc there is a detached shell of optical thin ionized gas which emits for free-free encounters. The ionized mass is $\\sim1.51\\,\\rm M_{\\odot}$ and the average density is $32\\,\\rm cm^{-3}$;\\\\ 5) almost co-spatial with the radio emission (Paper 1), there is a component emitting in the mid-IR, which is very likely due to line transitions (HeII, [NeII], [NeIII]) as found in this work. If a component of hot ($T\\sim340-550\\,\\rm K$) thermal dust is present, it has a negligible mass. We exclude the presence of non-thermal dust;\\\\ 6) a component of warm ($T\\sim 60-85\\,\\rm K$) grains is distributed in three concentric shells and has a total mass of $\\sim0.02\\,\\rm M_{\\odot}$ (in the case of mixed grains of silicate and graphite);\\\\ 8) no cool dust is evident in the far-IR images;\\\\ 9) the IR spectra show instead the presence of a PDR, very likely distributed outer than the ionized region, but the considered datasets do not allow us to evaluate the contribution of the interstellar material. The electron density in the second shell is yet lower than $\\sim 100\\,\\rm cm^{-3}$, with the exception of the outer region in the north-east, where the density is higher ($\\sim200\\,\\rm cm^{-3}$). This region corresponds to the radio ``spur-emission\", first detected by \\citet{1994Higgs} and present also in our radio maps . The spectral index analysis of this region gives evidence of the presence of density clumps as well. These clumps could be explained as region modified by the impact of the ejecta with the local environment;\\\\ 10) the gas-to-dust ratio in the nebula is $\\approx80 +\\rm \\frac{M_{neutral [M_{\\odot}]}}{0.02}$.  The nebula morphology suggests that the star has lost mass in at least three different episodes. As already reported in Paper 1, two of them occurred about  $5.4 \\times 10^{4}\\,\\rm yr$ and $2.7  \\times 10^{4}\\,\\rm yr$ ago \\citep[assuming an expansion velocity of the shell of $\\sim30\\,\\rm km\\,s^{-1}$,][]{1996Waters} and now we indicate an intermediate episode of $4.0 \\times 10^{4}\\,\\rm yr$. Considering these epochs, the current mass-loss does not explain the mass contained in the nebula. Assuming a linear expansion, this rate would have produced, in fact, a nebula of mass about few  $10^{-2}\\,\\rm M_{\\odot}$, i.e. two order of magnitude less massive than the value determined from the observations. This implies that a higher mass-loss rate took place to form such massive nebula and is consistent with changes expected in the rate of mass-loss for LBVs and observed in the case of AG Car.  \\citet{1996Waters} reported, by examining IRAS images, that the nebulae were ejected with a mass-loss rate of $ \\sim 5 \\times 10^{-4}\\,{\\rm M_{\\odot} yr^{-1}}$. By considering the nebular mass estimated in this work, we suggest instead a rate of $\\sim6.5\\times 10^{-5}\\,{\\rm M_{\\odot} yr^{-1}}$, which is consistent with many galactic LBVs. Mass-loss rates of the order of $ \\sim 10^{-5}\\,{\\rm M_{\\odot} yr^{-1}}$ have been derived for AG Car \\citep{2009Groh}, the Pistol star and FMM362 \\citep{2009Najarro}, G24.73+0.69 \\citep{2003Clark}, instead higher mass-loss (greater than $ \\sim 10^{-4}\\,{\\rm M_{\\odot} yr^{-1}}$) have been determined for the more luminous $\\eta$ Car \\citep{2001Hillier}, AFGL 2298 \\citep{2009Clark} and G26.47+0.02 \\citep{2012Umana}. Lower mass-loss rates ($ \\sim 10^{-6}\\,{\\rm M_{\\odot} yr^{-1}}$) were also observed in few cases, e.g. W243 \\citep{2009Ritchie}, Wra 17-96 \\citep{2002Egan} and HD162685 \\citep{2010Umana}.  In the computation of the total nebular mass we have not considered the contribution of the neutral gas. Therefore, our estimate of the nebular mass and the average mass-loss rate must be considered lower limits. However, the discrepancy with \\citet{1996Waters} is probably mostly due to a difference in the ionised mass determined in previous works with respect to our estimate. In fact, the mass determined in this work is at least four times smaller than the mass computed by \\citet{1994Higgs}. A massive nebula, such that proposed previously, would be comparable in mass to the Homunculus nebula around $\\eta$ Car, despite G79.29+0.46 is less luminous, and consequently less massive, than $\\eta$ Car. While the difference between the high nebular mass around G79.29+0.46 and its initial mass would have fascinating implications for massive stellar evolution, the ejecta mass estimated in this work appears to be in agreement with a 'low-luminosity' massive star.  G79.29+0.46 is still considered a candidate LBV because it does not match the variability criteria. However, indirect evidence for changes in the mass-loss rates have been pointed out. Indeed it could be a LBV in the quiescent state and which suffered at least three events of higher mass-loss over the last $5.4\\times10^{4}$ yr."
    },
    "1402/1402.0512_arXiv.txt": {
        "abstract": "The annihilation of dark matter particles  in the  halo of galaxies may end up into  $\\gamma$-rays, which travel almost unperturbed till to their detection at Earth. This annihilation signal can exhibit an anisotropic behavior quantified by the angular power spectrum, whose properties strongly depend on the dark matter  distribution and its clumpiness. We use high resolution pure dark matter N-body simulations to quantify the contribution of different components (main halo and satellites) to the global signal as a function of the analytical profile adopted to describe the numerical results. We find that the smooth main halo dominates the angular power spectrum of the $\\gamma$-ray signal up to quite large multipoles, where the sub-haloes anisotropy signal starts to emerge, but the transition multipole strongly depends on the assumed radial profile. The extrapolation down to radii not resolved by current numerical simulations can affect both the normalization $and$ the shape of the $\\gamma$-ray  angular power spectrum. For the sub-haloes described by an asymptotically cored dark matter distribution, the angular power spectrum shows an overall smaller normalization and a flattening at high multipoles. Our results show the criticality of the dark matter density profile shape in $\\gamma$-ray anisotropy searches, and evaluate quantitatively the intrinsic errors occurring when  extrapolating the dark matter radial profiles down to spatial scales not yet explored by numerical simulations. ",
        "introduction": "\\label{sec:intro}   One of the most reliable solutions to the missing mass in the Universe implies that it is constituted by  Weakly Interacting Massive Particles (WIMPs), clustered in galaxies as dark haloes. The astrophysical evidence of these particle dark matter (DM) candidates can be explored by direct as well as indirect detection techniques. In the latter case, the idea is that  WIMP DM may annihilate in pairs and produce charged particles and $\\gamma$-rays, detectable as rare components in cosmic rays. Differently from the charged cosmic rays, the flux of $\\gamma$-rays arriving at the Earth is not deflected by  magnetic fields and traces back directly to its sources. The search for DM through $\\gamma$-rays is therefore a preferential tool for exploiting the properties of its spatial distribution (\\cite{2012PDU.....1..194B} and refs. therein).  The excellent performances of the Large Area Telescope (LAT) on the  Fermi $\\gamma$-ray space Telescope ({\\it Fermi}) have let the exploration for a DM component in the Milky Way, in extragalactic nearby objects, as well as in cosmological structures \\citep{2012ApJ...761...91A,2010ApJ...712..147A,2010JCAP...05..025A,Abdo:2010dk, 2011PhRvL.107x1302A}. At high galactic latitudes, a faint $\\gamma$-ray irreducible emission has been measured, and shown to be isotropic at a high degree \\citep{Ackermann:2012uf}. The {\\it Fermi}-LAT has already reported the detection of a non-zero angular power spectrum (APS) above the noise level in the multipole range $\\ell \\sim 155\\div504$, corresponding to an angular scale $\\lsim 2^\\circ$, with a significance ranging from $5.3\\sigma$ between 2 and 5 GeV to $0.8\\sigma$ between 10.4 and 50 GeV \\citep{Ackermann:2012uf}.  Different predictions for the APS have been proposed for various populations of unresolved sources, both of astrophysical \\citep{2012PhRvD..86f3004C,SiegalGaskins:2010mp,2012JCAP...11..026H} and of DM origin \\citep{Fornasa09,2013MNRAS.429.1529F,2013PhRvD..87l3539A,SiegalGaskins:2008ge}. Since it is expected that the statistical properties of the DM distribution in galactic and extragalactic space are different from those of standard astrophysical objects, the study of the APS ascribable to DM sources may be an important signature worth to be explored.  The intensity of the $\\gamma$-ray signal depends on a particle physics term - describing the strength and the energy spectrum of the annihilation - and on the  DM density in collapsed structures. While the first factor includes the details of the assumed particle physics  model for the WIMPs, the second one has an astrophysical origin and it is usually modeled according to the results of cosmological collision-less simulations, generally predicting a steepening of the DM profile in the inner parts of the resolved halos (e.g. \\cite{aquarius,kuhlen08} and references therein).  Therefore, the most likely detectable targets have  been identified with especially dense regions such as the galactic center \\citep{Gomez-Vargas:2013bea,Fermi-LAT:2013uma}, and the center of any DM sub-structure orbiting in the Milky Way halo, like faint and ultra faint dwarf galaxies \\citep{Walker:2011fs,Ackermann:2013yva}.  The APS gives the measure of a signal correlation between two angular scales, and, in turns, between two spatial scales. For a source located at the galactic center, like the main halo of the Milky Way, the APS at multipoles, for example, $\\ell>500$ probes the DM distribution at $R < \\pi/500\\cdot8.5~\\kpc \\sim 40 ~\\pc$.  The study of the APS at $l\\gsim 500$ requires therefore to know the DM profile at scales much below the resolution ($\\sim$200 pc) of current state of the art numerical simulations for structure formation \\citep{aquarius,kuhlen08,2009MNRAS.398L..21S}. Several profile parameterizations provide excellent fits to the DM distribution of simulated halos \\citep{NFW,2004MNRAS.349.1039N,2006AJ....132.2701G,2008MNRAS.391.1940M,2009MNRAS.398L..21S}. However, when extrapolated below the resolution limit of cosmological simulations, different profiles predict very different central densities.   In this paper we discuss in detail these points when applied to the anisotropy in the $\\gamma$-ray flux from DM annihilation, namely: i) the intrinsic uncertainty due to the extrapolation to short distances of the DM distribution determined from numerical simulations; ii) the different signatures in the APS in connection with the various density profiles (cored and cuspy).  \\bigskip ",
        "conclusions": "We have calculated the intensity APS of the $\\gamma$-ray flux from DM annihilation in the halo of a Milky Way like galaxy, employing the original results from  recent numerical simulations of structure formation, which predict the DM phase space and its clustering into sub-haloes. The simulated galactic halo and its sub-haloes can be equally well  interpreted in terms of a peaked Einasto as well as an asymptotically cored MS radial dark matter profile.  We show here that the different parameterizations for the DM density distribution leads to very different predictions for the $\\gamma$-ray intensity APS. The DM halo and sub-haloes, when interpreted in terms of the peaked Einasto profile, yield much higher APS at small radial scales (high $l$) than the cored MS $\\rho(r)$.  The two profiles give comparable APS for the main halo only for $l\\lsim$ 10, while at $l$=100 the Einasto APS is about two orders of magnitude higher than the MS one. At very small scales, $l\\simeq$1000, the main halo within the MS profile does not contribute any longer to the anisotropy of the sky, while the Einasto profile still provides a sizable APS (about eight orders of magnitude above the MS contribution). We have also proven that the sub-haloes APS is significantly  lower for the MS case. Indeed, the APS of the sub-structures described by the Einasto profile is higher than the MS one since the sub-haloes, appearing as point-like sources, inject more angular power at almost all scales.  Our results demonstrate that the extrapolation of the radial DM profile down to radii not proven by cosmological simulations is specially dangerous when dealing with the search for anisotropies in the $\\gamma$-ray emission. The results for the APS at high multipoles may differ by huge amounts by a mere re-interpretation of the simulated haloes with a different DM radial density distribution. Also, depending of the assumed profile, it may occur that the sub-haloes give a peculiar signature in the APS or, at variance, that the main halo dominates at all multipoles the $\\gamma$-ray emission. In the latter case, the APS signature for DM annihilating in the galactic halo is significantly weakened.  As a final comment, we underline the caution in adopting extrapolated DM profiles when dealing with anisotropy searches, and emphasize the need for a better knowledge of the distribution of the DM in its clustered structures, especially taking into account the possible effects of baryonic matter (e.g. \\citet{Maccio2012, 2014MNRAS.437..415D}). On the other hand, $\\gamma$-ray anisotropy analysis will turn out to be crucial for probing the spatial DM distribution in the galaxy. Indeed, high multipoles measurements will probe scales well beyond the simulations' resolution and will help in discriminating the DM profile at very small radii."
    },
    "1402/1402.5930_arXiv.txt": {
        "abstract": "We present a study of sub-pc scale radio structure of the radio galaxy 3C~84/NGC~1275 based on the Very Long Baseline Array (VLBA) data at 43~GHz.  We discover a limb-brightening in the ``restarted\" jet associated with the 2005 radio outburst.  In the 1990s, the jet structure was ridge-brightening rather than limb-brightening, despite the observations being done with similar angular resolution.  This indicates that the transverse jet structure has changed recently.  This change in the morphology shows an interesting agreement with the $\\gamma$-ray flux increase, i.e., the $\\gamma$-ray flux in 1990s was at least seven times lower than the current one.  One plausible explanation for the limb-brightening is the velocity structure of the jet in the context of the stratified jet, which is a successful scenario to explain the $\\gamma$-ray emission in some active galactic nuclei (AGNs).  If this is the case, the change in apparent transverse structure might be caused by the change in the transverse velocity structure.  We argue the possibility that the transition from ridge-brightening to limb-brightening is related to the $\\gamma$-ray time variability on the timescale of decades.  We also discuss the collimation profile of the jet. ",
        "introduction": "The radio source 3C~84 is associated with the giant elliptical/radio galaxy NGC~1275 (z=0.0176).  Its classification by radio luminosity is Fanaroff-Riley type-I \\citep[e.g.,][]{Chiaberge1999}.  Thanks to its brightness and proximity, this source is one of the best-studied radio sources in history.  Recently increased activity starting in 2005 has been detected in radio band \\citep{Abdo2009}.  The Very Long Baseline Interferometry (VLBI) observations revealed that this flux density increase originated within the central pc-scale core, accompanying the ejection of a new jet component \\citep[hereafter Paper I]{Nagai2010}.  This new component appeared from the south of the core around 2003, and is moving to the position angle $\\sim160^{\\circ}$ steadily with slightly changing speed in both parallel and perpendicular directions \\citep[hereafter Paper II]{Suzuki2012}.  The apparent speed ranges from $0.1c$ to $0.5c$ (Paper II).  The moving direction of the new component is clearly different from the direction of the pre-existing component near the core separated about 0.5~mas along the position angle of $-150^{\\circ}$.  The flux density of both the new component and the core particularly increased after 2007-2008, and the new component has shown a further increase in flux density after 2009\\citep[hereafter Paper III]{Nagai2012}.  It is notable that 3C~84 is the best studied $\\gamma$-ray radio galaxy along with M~87 and Centaurus~A, and therefore these sources are the ideal laboratories to study the $\\gamma$-ray emission mechanism in misaligned Active Galactic Nuclei (AGNs).  Strong variability in the $\\gamma$-ray emission has been detected from 3C~84 by the Large Area Telescope (LAT) on board {\\em Fermi}.  An averaged $\\gamma$-ray flux during the first four months is $(2.10\\pm0.23)\\times10^{-7}$~ph cm$^{-2}$ s$^{-1}$ above 100~MeV.  This $\\gamma$-ray flux is seven times brighter than the upper limit estimated by EGRET/{\\em CGRO}.  It was claimed that the innermost jet of 3C~84 is the most likely source of $\\gamma$-ray emission because of the different $\\gamma$-ray activity level between EGRET era and {\\it Fermi} era \\citep{Abdo2009}.  During the first two years of observation, variations on the timescale of a month were observed by {\\em Fermi}-LAT in $\\gamma$-rays: one occurred in 2009 April-May \\citep{Kataoka2010}, and the other one occurred in 2010 June-August \\citep{BrownAdams2011}.  In particular, during the second flare photons up to 102.5 GeV were detected.  The largest increased activity in GeV band was reported on 2013 January 21 \\citep{Ciprini2013}.  An average daily $\\gamma$-ray flux is six times greater than the average flux reported in the second {\\em Fermi}-LAT catalog.  In Paper III, the radio time variability was studied to search for possible correlation with the $\\gamma$-ray variability, but no clear correlation was found on the timescale of $\\gamma$-ray time variation.  Neither new component ejection nor change in morphology associated with the $\\gamma$-ray flares were found by VLBI observations (Paper III).  The lack of significant changes in radio band for 3C~84 after the detection of high $\\gamma$-ray activity leaves the debate on the region responsible for the high-energy emission and its location still open.  Besides, the apparent speed detected by VLBI is relatively slower than the jet speed predicted from one-zone synchrotron-self Compton model or deceleration jet model unless the jet angle to the line of sight is very small ($<5^{\\circ}$: Paper II).  This indicates that the gamma-ray emitting region may have a higher Lorentz factor, and that the emission may be more strongly beamed than would be implied by the Lorentz factor estimated from the VLBI proper motions.  We also performed spectral energy distribution (SED) fit to the observed broadband spectrum using the estimated apparent speed, but it failed to reproduce the optically thin radio spectrum observed by VLBI(Paper II).  While several possibilities to reconcile the discrepancy between radio and $\\gamma$-ray properties have been discussed in earlier works (Paper I; Paper II; Paper III), it is of great interest to investigate the presence of stratified structure with a velocity gradient \\citep{Ghisellini2005} in the jet of 3C~84.  % According to this scenario, the radio emission is mostly coming from the slower sheath while the emission from the spine is beaming away from the line of sight.   Therefore, the limb-brightening can be observed along the jet if the spine-sheath structure is present.  Clear evidence of limb-brightening is found in several AGNs such as M~87\\citep{Junor1999}, Mrk~501\\citep{Giroletti2008}, Mrk~421\\citep{Piner2010}, and 1144+35\\citep{Giovannini1999}, but no clear signature of limb-brightening has been found in 3C~84 so far.  To detect the limb-brightening, high spatial resolutions that can resolve the transverse direction of the jet and high dynamic range image are required.  In this paper, we report a new VLBA observation at 43~GHz to investigate the transverse structure of the jet in 3C~84.  Mainly we focus on discussing the origin of limb-brightening, but we also argue the collimation profile that can be obtained by resolving transverse direction.  At the 3C 84 distance 1~mas corresponds to 0.344~pc assuming $H_{0}$=70.5~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$, and $\\Omega_{\\Lambda}=0.73$.  Assuming a super massive black hole (SMBH) mass of $3.4\\times10^{8}~M_{\\bigodot}$ \\citep{Wilman2005}, $10^{3}$ Schwarzschild radii ($r_{g}$) corresponds $9.5\\times10^{-2}$~mas. ",
        "conclusions": "Clear limb-brightening is discovered in the sub-pc scale jet of 3C~84, which has been formed by the restarted activity observed since 2005, while the jet structure was rather ridge-brightening in the 1990s.  Although several possibilities are considered, we propose that the change from ridge-brightening to limb-brightening may be attributed to a change in the transverse velocity structure on the basis of a ``spine-sheath\" scenario.  This is compatible with the lack of the radio counterpart of short-term $\\gamma$-ray flares reported in Paper II.  The jet width profile in $\\sim10^{3}$-$10^{5}$~$r_{g}$ scale is rapidly collimated rather than parabolic shape and differs from the trend in M~87, which might reflect the different circumnuclear environment between M~87 and 3C~84.  Higher resolution studies are important to probe whether the limb-brightening is present in the innermost region of 3C~84 and study the collimation profile as well.  Future high-resolution instruments, such as (sub)mm VLBI including phase-up ALMA, will give us a hint for the jet physics in connection with the $\\gamma$-ray emission and jet formation theory."
    },
    "1402/1402.5101_arXiv.txt": {
        "abstract": "{Current and future imaging surveys will measure cosmic shear with a statistical precision that demands a deeper understanding of potential systematic biases in galaxy shape measurements than has been achieved to date. We investigate the effects of using the point spread function (PSF) measured with stars to determine the shape of a galaxy that has a different spectral energy distribution (SED) than the star. We demonstrate that a wavelength dependent PSF size, for example as may originate from atmospheric seeing or the diffraction limit of the primary aperture, can introduce significant shape measurement biases. This analysis shows that even small wavelength dependencies in the PSF may introduce biases, and hence that achieving the ultimate precision for weak lensing from current and future imaging surveys will require a detailed understanding of the wavelength dependence of the PSF from {\\em all} sources, including the CCD sensors. } ",
        "introduction": "\\label{sec:intro}  A goal of large astronomical imaging surveys is to constrain cosmological parameters by measuring the small departure from statistical isotropy of the shapes and orientations of distant galaxies, induced by the gravitational lensing from foreground large-scale structure. The shapes of galaxy images, however, are not only affected by cosmic shear (typically a $\\lesssim 1\\%$ shift in the major-to-minor axis ratio), but are also determined by the combined point spread function (PSF) due to the atmosphere (for ground-based instruments), telescope optics, and the image sensor -- together often a few $\\%$ shift. The size and shape of this additional convolution kernel is typically determined from the observed images of stars, which are effectively point sources before being smeared by the PSF. Galaxy images can then be deconvolved with the estimated convolution kernel. Implicit in this approach is the assumption that the kernel for galaxies is the same as the kernel for stars. If the PSF is dependent on wavelength, this assumption is violated since stars and galaxies have different spectral energy distributions (SEDs) and hence different PSFs. Correcting observed images with the incorrect PSF can lead to biases in shape measurements. To predict and eliminate these biases, we must understand all PSF contributions that depend on wavelength.  In this paper, we illustrate the biases that can arise from a wavelength-dependent PSF by considering a particular class of chromatic effects. In Section \\ref{sec:size}, we describe our PSF model, and in Section \\ref{sec:biases} estimate the resulting galaxy shape biases. In Section \\ref{sec:requirements}, we investigate how these biases affect current and proposed weak lensing surveys. We conclude in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Comparing Figures \\ref{fig:f1} and \\ref{fig:f2} to the rightmost column of in Table \\ref{tab:requirements}, we see that the mismatch in PSF size between stars and galaxies from effects such as chromatic seeing for a ground-based telescope and the diffraction limit for a space-based telescope is significantly larger than the requirements for cosmic shear. Fortunately, corrections can be applied on an object-by-object basis given we have some estimate of each object's SED over the wavelength range of the shape measurement filter(s). Such estimates are readily available through multifilter photometry. This approach, roughly akin to photometric redshifts, is studied in more detail in \\cite{Cypriano++10, Plazas+Bernstein12, Meyers+Burchat14}.  Corrections of this type require knowledge of the chromatic effects from the entire imaging system, including the atmosphere for ground-based telescopes, and of optics and sensors for both ground and space telescopes. For the class of chromatic effects investigated here ($\\mathrm{FWHM} \\propto \\lambda^\\alpha$), we can estimate how well we need to know the power law index $\\alpha$ in order for our correction to succeed. From Figure \\ref{fig:f1} and Table \\ref{tab:requirements}, we see that, for LSST $r$-band images, $\\delta r^2_\\mathrm{PSF}/r^2_\\mathrm{PSF}$ varies by a factor of $\\sim 25$ more than the requirement. Since $\\delta r^2_\\mathrm{PSF}/r^2_\\mathrm{PSF}$ varies roughly linearly with $\\alpha$, we therefore need to know $\\alpha$ to a precision of about $|\\alpha|/25 \\approx 0.008$. Similarly, knowledge of $\\alpha$ also needs to be at the level of $\\approx 0.01$ for Euclid. Of course, not all chromatic effects are accurately described by a power law model (e.g. \\cite{Plazas+Bernstein12}), but this exercise demonstrates the unprecedented requirements for understanding all aspects of future imaging systems."
    },
    "1402/1402.0724_arXiv.txt": {
        "abstract": "Recent studies about collisionless shocks in partially ionized plasmas showed that some of neutral particles leak into the shock upstream region from the downstream region. In this paper, we perform a linear analysis and show that acoustic waves are unstable in the neutral precursor region. The acoustic instability amplifies fluctuations of magnetic field and density in the upstream region. The fluctuations are indispensable for the diffusive shock acceleration and could be important for the downstream turbulence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1514_arXiv.txt": {
        "abstract": "We determine the extinction law through Cep OB3b, a young cluster of 3000 stars undergoing gas dispersal.  The extinction is measured toward 76 background K giants identified with MMT/Hectospec spectra. Color excess ratios were determined toward each of the giants using $V$ and $R$ photometry from the literature,  $g$,$r$,$i$ and $z$ photometry from SDSS and $J$, $H$, and $K_{s}$  photometry from 2MASS.  These color excess ratios were the used to construct the extinction law through the dusty material associated with Cep OB3b.  The extinction law through Cep OB3b is intermediate between the $R_{V} = 3.1$ and $R_{V} = 5$ laws commonly used for the diffuse atomic ISM and dense molecular clouds, respectively.  The dependence of the extinction law on line-of-sight $A_{V}$ is investigated and we find the extinction law becomes shallower for regions with  $A_{V} > 2.5$ magnitudes.  We speculate that the intermediate dust law results from dust processing during the dispersal of the molecular cloud by the cluster ",
        "introduction": "Young star clusters, with ages less than 10 Myr, are associated with dusty molecular clouds that are denser and colder than the diffuse interstellar medium (ISM).  The high density and cold temperatures of a molecular cloud allow dust grains to grow through mantle growth and agglomeration \\citep{ossenkopf1993,weid1994, ormel2009}.  These processes change the size distribution of the dust grains and thus the amount of attenuation of stellar light along a line-of-sight through the dust.  The extinction law for a line-of-sight measures this attenuation as a function of wavelength and is found to vary among lines of sight throughout the galaxy \\citep{ccm89,rl85}.  The extinction law can be parameterized by the ratio of total to selective extinction, given as \\begin{equation} R_{V}=\\frac{A_{V}}{E(B-V)} \\end{equation} Where $A_{V}$ is the total $V$-band extinction and \\begin{equation} E(B-V) = ((B-V)_{observed} - (B-V)_{intrinsic}) \\end{equation} is the $B-V$ color excess.  For large dust grains, the extinction becomes increasingly gray, that is, there is reduced wavelength dependence.  Thus, the selective extinction is reduced and $R_{V}$ is increased.  The diffuse ISM typically has $R_{V}=3.1$ \\citep{rl85}, whereas, in dense molecular clouds $R_{V} \\sim 5.6$ is typical \\citep{ccm89}.  Understanding the extinction law in molecular clouds is essential for studies of star formation for several reasons.  First, the slope of the extinction law can shed light on the properties of the dust and dust agglomeration in molecular clouds.  Furthermore, by changing the opacity, dust grains can affect important feedback processes such as the ionization of the cloud \\citep{weid1994}.  Finally, properly measuring the extinction toward young stars is critical for measuring the luminosity of the young stars in order to construct an HR diagram.   Located at a distance of 700~pc, as given by Very Long Baseline Array parallaxes to the neighboring Cep A region \\citep{moscadelli2009, dzib2011}, Cep OB3b is a young OB association on the North-West edge of the Cepheus molecular cloud.  Fig.~\\ref{fig_extdistro} shows the Cep OB3b region of the Cepheus molecular cloud.  Associated with the O and B stars is one of the largest young clusters within 1 kpc of the Sun.  ${\\it Spitzer}$ has revealed roughly 1000 young stars with infrared excesses indicative of young protoplanetary disks.  ${\\it Chandra}$ X-ray observations reveal that $2/3$ of the young stars have already dispersed their protoplanetary disks.  Based on these observations, \\citet{allen2012} estimate there are ~3000 members in this young cluster.  Age estimates in the literature for this young cluster are on the order of ~3-6 Myr \\citep{mayne2007,littlefair2010, bell2013}.  The Cep OB3b young cluster is composed of two main sub-clusters, the eastern of which has a lower peak stellar density and a lower disk fraction than the western, consistent with an older average age for the eastern sub-cluster.  Situated on the edge of the Cepheus Molecular cloud, Cep OB3b straddles the interface between the diffuse atomic material of the ISM and the dense molecular material of the cloud.  Although associated with the dense molecular clumps Cep B and Cep F \\citep{sargent77}, most of the young stars with protoplanetary disks are located in a cavity with an internal extinction $A_{V} < 5$ magnitudes \\citep{allen2012}.  This is consistent with a young cluster that has recently dispersed much of its natal molecular material.  Further, unlike the closer Orion Nebula Cluster, which is located in front of ionized nebular material, Cep OB3b is primarily located adjacent to ionized nebular material.  This configuration is advantageous because it allows measurements of background stars through the recently evacuated cavity as well as through the denser material closer to the molecular cloud.   We determine the extinction law through Cep OB3b by measuring the color excesses of background K giants in the optical through near-infrared regimes.  We find that the slope of the extinction law becomes shallower in regions of high extinction, indicative of dust grain growth.  In $\\S$~\\ref{sec:data} we describe the optical and infrared data used in this study.  In $\\S$~\\ref{sec:sources} we describe the identification of the background giant stars with MMT/Hectospec spectroscopy.  In $\\S$~\\ref{sec:extlaw} we derive the extinction law through Cep OB3b.  Finally, in $\\S$~\\ref{sec:discuss} we examine how the extinction law varies with line-of-sight extinction and suggest that the extinction law in the cavity is intermediate between ISM and GMC extinction laws. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Estimation of Foreground Extinction}  We estimate the amount of extinction along the line-of-sight due of the cluster by calculating the $J - K$ color excess to the diskless (class III) members identified by ${\\it Chandra}$ in \\citet{allen2012}.  Fig.~\\ref{fig_av_hist} shows the distribution of line-of-sight extinction estimates for the class III cluster members.  We assume the young stars have the intrinsic colors of  \\citet{luhman2010} and calculate the line-of-sight extinction for each star.  Also shown in fig.~\\ref{fig_av_hist} is the distribution of the line-of-sight extinction estimates for the background giant stars assuming they have the intrinsic colors of \\citet{ducati2001}.  Both distributions shows a sharp increase between $1- 2~A_{V}$; we interpret this increase as due to the dust in the cluster.  Therefore, we estimate there is roughly 1 $A_{V}$ of foreground extinction.    \\subsection{An evolving extinction law in the interface between the diffuse and molecular ISM?}  Previous studies of star forming regions have examined the change in the dust properties in dense molecular clouds with respect to the diffuse ISM. Recent works have found variations in the extinction laws as a function of  $A_{V}$ within a given molecular cloud \\citep{whittet2001, campeggio2007, moore2005}.  \\citet{whittet2001} examined the optical and near infrared photometry of individual field stars that are background to the Taurus Dark Cloud and calculated  $A_{V}$ and $R_{V}$ for each star. By correlating the optical depth of the 3 $\\mu$m ice absorption feature with the estimated  $A_{V}$ for each star, they determined an $A_{V}$ threshold for ice extinction in the Taurus Dark Cloud  of  $A_{V} \\sim 3.2$. Below this threshold, the value of $R_{V}$ is similar to the ISM, while above this threshold $R_{V}$ is higher. This suggests a sharp change in the dust properties at the threshold $A_{V}$, possibly corresponding to the growth of icy mantles on the dust grains.  \\citet{campeggio2007} use optical and near infrared photometry to measure changes in the extinction law through the dark globule CB 107. They find that $R_{V}$ increases with increasing  $A_{V}$, but instead of finding a threshold  $A_{V}$ value, this increase occurs throughout the range of observed extinctions. This finding is interpreted as evidence that coagulation is the primary mechanism for dust growth in CB 107 as opposed to ice mantle condensation.  The change in the dust law is thought to result from the growth of the grains.  Since the condensations of ices on grains can only lead to a relatively small change in the dust grain diameter \\citep{draine1985}, it is thought that most of the grain growth is through the coagulation of grains.  Theoretically, models show that coagulation of icy grains in dense molecular cloud environments will lead to an evolving dust law with time \\citep{ormel2009,ormel2011}.  The Cep OB3b cluster exists at the interface of a molecular cloud  and the diffuse ISM created by the dispersal of the cloud material by the young stars in the cluster.  Within this interface, we ave found an extinction law that is intermediate to that of the diffuse ISM and molecular cloud ISM.  We also find that the extinction law shows evidence of variations between the cluster and the molecular cloud region.  Previously, International Ultraviolet Explorer (IUE) observations by \\citet{massa1984} examined the ultraviolet (UV) extinction curves toward O and B stars in the Cep OB3 association. They found a correlation between the value of the far-UV extinction, parameterized by $\\frac{E(15-V)}{E(B-V)}$, where 15 refers to the IUE bandpass centered at 1550 \\AA, and distance from the molecular cloud. Stars near the molecular material have low $\\frac{E(15-V)}{E(B-V)}$ whereas stars more distant from the molecular material have $\\frac{E(15-V)}{E(B-V)}$ values close to the galactic mean. Because the UV extinction depends strongly on the size distribution of the dust grains, \\citet{massa1984} suggested that the dust properties change with distance from the cloud. Lines of sight distant from the cloud have smaller mean grain sizes than lines of sight near the cloud.  We propose that the transition between diffuse ISM and molecular cloud dust laws is not sharp for molecular clouds associated with large clusters, or clusters that carve a cavity in their molecular material.  In these regions, cloud material may be returned to the ISM when it is disrupted by outflows and winds or if it is carried off the cloud by the flows of material photoevaporated off the cloud surfaces by the UV radiation from the O-stars.  Indeed, the 24~$\\mu$m emission apparent inside the cavity in the Cep OB3b is strongest just outside the cloud surface being evaporated by the O-star HD~217086 (Fig.~1, \\citet{allen2012}); this emission appears to be dust carried away in the flow hot, ionized, photoevporated gas. Once removed from the molecular cloud environments, a variety of feedback processes can affect dust grains, such as heating by radiation, which can evaporate ice mantles, and collisions in shocks driven by stellar winds, outflows and ionization fronts which can break up agglomerated dust grains. In this explanation, the intermediate extinction law apparent in the Cep OB3b cloud is due to processing of dust grains from the molecular cloud, which over the few million year lifetime of the cluster, have not yet returned the dust back to the small grain population typical for the diffuse ISM."
    },
    "1402/1402.4794_arXiv.txt": {
        "abstract": "Some models of coronal heating suppose that convective motions at the photosphere shuffle the footpoints of coronal magnetic fields and thereby inject sufficient magnetic energy upward to account for observed coronal and chromospheric energy losses in active regions. Using high-resolution observations of plage magnetic fields made with the Solar Optical Telescope aboard the {\\em Hinode} satellite, we investigate this idea by estimating the upward transport of magnetic energy --- the vertical Poynting flux, $S_z$ --- across the photosphere in a plage region.  To do so, we combine: (i) estimates of photospheric horizontal velocities, $\\vvec_h$, determined by local correlation tracking applied to a sequence of line-of-sight magnetic field maps from the Narrowband Filter Imager, with (ii) a vector magnetic field measurement from the SpectroPolarimeter.  Plage fields are ideal observational targets for estimating energy injection by convection, because they are: (i) strong enough to be measured with relatively small uncertainties; (ii) not so strong that convection is heavily suppressed (as within umbrae); and (iii) unipolar, so $S_z$ in plage is not influenced by mixed-polarity processes (e.g., flux emergence) unrelated to heating in stable, active-region fields. In this plage region, we found that the average $S_z$ varied in space, but was positive (upward) and sufficient to explain coronal heating, with values near $(5 \\pm 1) \\times 10^7$ erg cm$^{-2}$ s$^{-1}$. We find the energy input per unit magnetic flux to be on the order of $10^5$ erg s$^{-1}$ Mx$^{-1}$.  A comparison of intensity in a Ca II image co-registered with one plage magnetogram shows stronger spatial correlations with both total field strength and unsigned vertical field, $|B_z|,$ than either $S_z$ or horizontal flux density, $B_h$. The observed Ca II brightness enhancement, however, probably contains a strong contribution from a near-photosphere hot-wall effect, which is unrelated to heating in the solar atmosphere. ",
        "introduction": "How is the solar corona heated to temperatures of $\\sim 1$ MK, when the lower layers of the solar atmosphere are $\\sim 10^4$ K or less? Evidently, the energy needed to heat the Sun's atmosphere must cross the photosphere in some organized form before being converted into heat (disorganized, ``thermalized'' energy) in the chromosphere and corona.  Because the magnetic fields that permeate the corona are all anchored at the photosphere, they are natural candidates for energetic coupling between the solar interior and corona.  In the interior, motions in the Sun's gas are driven by convection, and some fraction of the kinetic energy in turbulent convective motions is thought to be converted into energy stored in electric currents flowing in coronal magnetic fields that is then dissipated as heat. These induced currents might be characterized as either steady or rapidly varying (e.g., wave-driven) relative to the timescales of the atmospheric response, and the dissipation of each has been referred to as ``DC'' (direct-current) or ``AC'' (alternating-current) heating, respectively (e.g., Klimchuk 2006). \\nocite{Klimchuk2006}  To be a viable coronal heating mechanism, the input energy must be commensurate with observed energy losses in active region (AR) fields, estimated by \\citet{Withbroe1977} to be $\\sim 1 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$ for the corona and $\\sim 2 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$ for the chromosphere.  Waves were once thought to be primarily responsible for coronal and chromospheric heating (see, e.g., Withbroe and Noyes 1977).  \\nocite{Withbroe1977} While waves (e.g., Tomczyk \\etal \\,2007) \\nocite{Tomczyk2007} and wave dissipation (e.g., Hahn \\etal \\,2012) \\nocite{Hahn2012} have been reported in the corona, currently available observations have not demonstrated that they supply sufficient energy to heat the active-region chromosphere, transition region, and corona.  In contrast to models invoking dissipation of waves, other models posit that convective motions induce relatively long-lived, DC currents that are episodically dissipated to heat the chromosphere and corona.  We explore the latter hypothesis here.  There is a long history of modeling this convection-driven coronal energy input.  Parker (1983a, 1983b) \\nocite{Parker1983, Parker1983b} proposed that convection braids and twists the photospheric footpoints of coronal magnetic fields, and thereby injects energy into the corona.  This energy is stored in current sheets, and is transiently dissipated in small bursts referred to as nanoflares \\citep{Parker1988}, with typical energies of $\\sim 10^{24}$ erg, about $10^{-9}$ of the energies in very large flares. \\citet{Galsgaard1996} modeled an idealization of this process by imposing shearing flows on the upper and lower boundaries of an initially uniform field in an MHD simulation, and found sufficient power to heat the corona. \\citet{Gudiksen2002} imposed a more complex flow field, meant to mimic convective motions, on an MHD model of the coronal field and also found sufficient power, as well as morphology consistent with aspects of coronal observations.  In the framework of reduced MHD, \\citet{Rappazzo2008}, also found sufficient power, even though fields in their model were only weakly braided. More recently, \\citet{Bingert2011} also modeled this process in MHD with a detailed treatment of the energy equation and found heating that is transient in time and space, and concentrated in and near the modeled transition region.  One promising observational approach to constraining models of coronal heating is to analyze time evolution of magnetic fields at the photosphere, where the magnetic field is precisely and routinely measured.  Clear evidence of braiding or twisting motions would support the mechanism proposed by Parker.  \\citet{Schrijver1998} proposed that continuous emergence and cancellation of small-scale fields in the quiet Sun's ``magnetic carpet'' leads to reconnection and heating, but \\citet{Close2004} used sequential potential models of quiet-Sun fields to argue that emergence and cancellation are not required: just reconnection between existing flux systems, as their photospheric footpoints move, should be sufficient.  \\citet{Meyer2013} recently directly incorporated magnetogram sequences into the lower boundary of a magnetofrictional model of quiet-sun coronal field evolution, to investigate the dissipation of magnetic energy within the simulation.  Aspects of energy dissipation in their model were qualitatively consistent with solar observations, although their total upward energy flux was smaller than the observationally estimated energy demand for the quiet-sun atmosphere.  \\citet{Yeates2014} recently investigated analytic expressions for lower bounds on the upward-directed Poynting flux of magnetic energy in a region of plage fields in NOAA AR 10930, based upon observed photospheric magnetic and velocity fields.  The flows they analyzed were estimated by \\citet{Welsch2012}, who applied Fourier Local Correlation Tracking (FLCT; Fisher and Welsch 2008) \\nocite{Fisher2008} to a sequence of line-of-sight magnetograms (magnetic field maps) of this active region.  These magnetograms were observed with the Narrowband Filter Imager (NFI) instrument on the Solar Optical Telescope (SOT) \\citep{Tsuneta2008, Suematsu2008, Ichimoto2008, Shimizu2008} aboard the {\\em Hinode} satellite \\citep{Kosugi2007}, with a cadence $\\simeq 120$ s, over about 13 hours on 2006 December 12 -- 13.  \\citet{Yeates2014} compared their lower bounds on the Poynting flux with a direct estimate of the Poynting flux, obtained with a procedure that we explain in detail here. First, they assumed that the photospheric magnetic field, $\\bvec$, is frozen to the plasma --- a valid assumption in quite general circumstances (see, e.g., Parker 1984). \\nocite{Parker1984a} Then the photospheric electric field, $\\evec$, is ideal, and equal to $-(\\vvec \\times \\bvec)/c$, where $\\vvec$ is the photospheric velocity.  Then the (vector) Poynting flux of magnetic energy, $\\Svec,$ can be expressed in terms of $\\vvec$ and $\\bvec$ as \\be \\Svec = \\bvec \\times (\\vvec \\times \\bvec)/4 \\pi ~. \\label{eqn:poynting_vec} \\ee Approximating the photospheric surface as locally planar, we adopt Cartesian geometry, and use $z$ and $h$ to refer to vertical and horizontal directions, respectively.  Then the vertical component of the Poynting flux is \\be S_z = [v_z B_h^2 - (\\vvec_h \\cdot \\bvec_h)B_z]/4\\pi ~. \\label{eqn:poynting_vert} \\ee This expression for total Poynting flux has been conceptually divided into an ``emergence'' term, which contains $v_z$, and a ``shear'' term, which contains $\\vvec_h$ \\citep{Liu2012d, Parnell2012}.  We digress for a moment to note that a positive (upward) value for the shearing term also implies the emergence of magnetized plasma across the photosphere. A clear example of this is the special case in which $v_z$ is zero (so the emergence term vanishes), and the shearing term is positive. Then both $B_z$ and $\\bvec_h$ must be nonzero, implying $\\bvec$ is tilted; and $\\vvec_h$ must have a nonzero projection onto $\\bvec_h$. We define the component of $\\vvec_h$ along the horizontal field $\\bvec_h$ to be $\\vvec_{Bh}$, and then further decompose $\\vvec_{Bh}$ into a component parallel to the total field $\\bvec$, which we label $v_{Bh,\\parallel}$, and a component perpendicular to $\\bvec$, which we label $v_{Bh,\\perp}$. Since $\\bvec$ is tilted, $v_{Bh,\\perp}$ must also be; and it must be tilted upward when the shearing term is positive. This upward tilt for $v_{Bh,\\perp}$ implies that this component of the flow advects the tilted magnetic flux upward. The parallel flow $v_{Bh,\\parallel}$ neither advects magnetic field or produces a Poynting flux. The counter-intuitive result that a horizontal velocity can produce upward transport of magnetic fields arises because the total velocity included a component of the velocity parallel to $\\bvec$, which is irrelevant for the Poynting flux. (In the special case that the parallel velocity is zero, then equation [\\ref{eqn:poynting_vec}] reduces to $\\Svec = B^2 \\vvec/4\\pi$, and equation [\\ref{eqn:poynting_vec}] becomes $S_z = B^2 v_z/4\\pi$.)  Since \\citet{Yeates2014} were primarily focused on heating in plage --- regions of nearly-vertical field when new flux is not emerging --- the shearing term should dominate, meaning \\be S_z^{\\rm \\,plage} \\simeq - (\\vvec_h \\cdot \\bvec_h)B_z/4\\pi ~. \\label{eqn:poynting_shear} \\ee  \\citet{Yeates2014} treated the flows estimated by FLCT as horizontal velocities.  We note that there is some controversy about how to interpret of velocities determined by correlation tracking and other ``optical flow'' \\citep{Schuck2006} methods.  \\citet{Demoulin2003} suggested that the apparently horizontal flows estimated by LCT are a linear combination of the horizontal velocity with the vertical velocity, with weighting determined by the ratio of horizontal to vertical magnetic field. To test the accuracy of velocities reconstructed from magnetogram sequences, \\citet{Welsch2007} compared flows estimated by several methods, including LCT, using synthetic magnetograms extracted from MHD simulations of an emerging magnetic flux tube in the solar interior in which the actual velocities were known. Using the same test data, \\citet{Schuck2008} subsequently argued that optical flow methods, such as LCT, essentially estimate the horizontal velocity, $\\vvec_h$, although their estimates can be affected by vertical flows.  The NFI magnetograms only provide estimates of the line-of-sight (LOS) field, $B_{\\rm LOS}$, but the expressions for the Poynting flux given above all require knowledge of the vector magnetic field, $\\bvec$. Accordingly, \\citet{Yeates2014} co-registered the (12 Mm $\\times$ 12 Mm) region of the NFI field of view (FOV) that they studied with the corresponding sub-region of a vector magnetogram observed by SOT's SpectroPolarimeter (SP; Lites \\etal \\, 2013). \\nocite{Lites2013} A reprojected vector magnetogram based upon these observations was prepared by \\citet{Schrijver2008} and is available online. The co-alignment procedure followed the approach used by \\citet{Welsch2012}, described in their Appendix.  By combining $\\vvec_h$ estimated with FLCT with $\\bvec$ from the SP magnetogram, \\citet{Yeates2014} estimated the average Poynting flux to be $1.7 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$.  This energy flux is less than the combined energy demand for the chromosphere and corona in active regions estimated by \\citet{Withbroe1977}. As discussed in greater detail below, however, this estimate did not account for the observationally estimated magnetic filling factors that had been applied to each magnetic field component in the vector magnetogram used by \\citet{Schrijver2008}. Insufficient Poynting flux would indicate that processes on spatial or temporal scales that are unresolved in these photospheric observations (e.g., waves or smaller-scale footpoint shuffling) play a significant role in heating.  Despite the central role of the Poynting flux in theories of coronal heating, very few observational estimates of Poynting flux in the context of coronal heating have been published. \\citet{Tan2007} investigated a ``proxy Poynting flux,'' $|\\vvec_h| B_{\\rm LOS}^2$ in more than 160 active regions, determined by applying LCT to LOS magnetograms, and found values in the range $10^{6.7}$ -- $10^{7.6}$ erg cm$^{-2}$ s$^{-1}$. They used the LOS fields alone because sequences of vector magnetic field measurements were quite rare.  More recently, \\citet{Kano2014} estimated the work done by flows on the magnetic field, by parametrizing the expected deformation of coronal loop structure by surface flows.  The expression they derive is inversely proportional to coronal loop length $L \\sim 100$ Mm, and for LCT flows with typical magnitudes of 0.5 km s$^{-1}$, they estimate a Poynting flux of $2 \\times 10^6$ erg s$^{\u22121}$ cm$^{\u22122}$. They also found typical flow speeds were lower in regions with higher magnetic filling factors.  This is roughly an order of magnitude lower than the value of $1.7 \\times 10^7$ erg s$^{\u22121}$ cm$^{\u22122}$ reported by \\citet{Yeates2014}.  Given this considerable variation in published Poynting flux estimates, further investigation of Poynting fluxes is warranted.  Here, we report additional estimates of the Poynting flux from plage magnetic fields in the same active region studied by \\citet{Yeates2014}.  Our primary aim is to investigate the properties of photospheric Poynting flux in greater detail than was done previously, including its dependence on photospheric magnetic field structure. The remainder of this paper is organized as follows.  In the next section, we briefly describe the magnetic field data and tracking methods we used to estimate $\\vvec_h$. In Section \\ref{sec:results}, we first present our estimates of the Poynting flux in another small region of plage in AR 10930, then analyze the Poynting flux's correlations with magnetic structure in the region. The region of the NFI FOV that we analyze here was also observed in Ca II by SOT's Broadband Filter Imager (BFI; Tsuneta \\etal \\,2008), \\nocite{Tsuneta2008} and in \\S \\ref{subsec:ca_ii} we compare this chromospheric emission with the spatial distributions of Poynting flux and magnetic field components. Finally, we conclude with a brief discussion of our results in Section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  By combining LCT velocities estimated from a LOS magnetogram sequence with a vector magnetogram, both derived from {\\em Hinode}/SOT observations of AR 10930, we estimated the Poynting flux, under the frozen-in-flux assumption, in a (12 Mm $\\times$ 12 Mm) plage region to be 4.9 -- $5.5 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$, depending upon whether FLCT or SLCT velocities were used.  Errors in the magnetic fields likely produce smaller uncertainties in the Poynting flux than this.  These Poynting fluxes are greater than the chromospheric and coronal energy demands estimated by \\citet{Withbroe1977}, $\\sim 2 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$ and $\\sim 1 \\times 10^7$ erg cm$^{-2}$ s$^{-1}$, respectively.  We found that the Poynting flux varied in sign across the plage region that we studied.  By plotting the distribution of Poynting fluxes in this region's pixels, we found that the mean upward flux arose from a predominance of upward-flux pixels toward the high-Poynting-flux end of the distribution.  We then identified ``plage-like'' pixels --- those with nearly vertical flux densities in the range 100 -- 1500 Mx cm$^{-2}$, and sufficient polarization for the vector field to be estimated --- across the rest of the active region.  The distribution of Poynting fluxes in this set of plage-like pixels exhibited the same systematic prevalence of upward-flux pixels, suggesting that the plage region that is the focus of our study is not a special case.  In analogy with the study by \\citet{Fisher1998} relating active regions' soft X-ray luminosities to magnetic field properties, we investigated correlations of Poynting fluxes with properties of the magnetic and velocity fields.  We found that both the {\\em unsigned} and {\\em signed} Poynting fluxes typically increase with pixels' field strengths.  Correlations of Poynting fluxes with both unsigned vertical electric current density and flow vorticity were relatively weak, suggesting braiding or vortical motions (e.g., Parker 1983, Kitiashvili et al. 2014) \\nocite{Parker1983, Kitiashvili2014} are not key aspects of the energy transport process.  Building upon the work of \\citet{Fisher1998}, \\citet{Pevtsov2003} found that soft X-ray luminosities for a range of magnetic regions on the Sun scaled nearly linearly with flux, with a relationship approximating $\\sim 10^3$ erg s$^{-1}$ Mx$^{-1}$.  Here, we found the energy input per unit magnetic flux to be on the order of $10^5$ erg s$^{-1}$ Mx$^{-1}$.  We found that fields with intrinsic vertical field strengths of $\\sim 1300$ G supply the bulk of the net Poynting flux.  We also compared our Poynting flux map with a Ca II intensity image, and found much stronger correlation of Ca II emission with the vertical magnetic field strength $|B_z|$ than with vertical Poynting flux.  We noted that this magnetic correlation might, however, arise from near-photospheric emission in the passband exhibiting the hot-wall effect in strong fields.  The time interval $\\Delta t$ between images (eight minutes here) and windowing length scale $L$ (4 pixels, $\\sim 1$ Mm) used in our tracking will likely filter out processes on shorter temporal and spatial scales.  Such processes (e.g., waves, or smaller-scale braiding) might play key roles in chromospheric emission.  Given that chromospheric and coronal length scales are shorter than the scales we resolve, observations with higher resolution in space and time (see below) would be useful to investigate Poynting flux -- emission correlations further.  It should be noted, however, that our energy flux is large enough that sub-resolution dynamics are not required to explain the observed coronal heating.  This initial study leaves several questions unanswered, motivating related studies to extend the work here.  Do other tracking methods yield similar results?  The same plage region analyzed here could be tracked with other methods (e.g., DAVE or DAVE4VM; Schuck 2006, Schuck 2008) \\nocite{Schuck2006, Schuck2008} to better understand the model-dependence of flow estimates in determining Poynting fluxes. How are photospheric Poynting fluxes related to emission from the overlying atmosphere?  To address this question, it would be useful to analyze additional {\\em Hinode}/SOT datasets, especially observations with simultaneous IRIS \\citep{DePontieu2014} coverage of chromospheric, transition region, and coronal emission, to seek any evidence of spatial or temporal correlations between energy input via our estimated Poynting fluxes and energy dissipation in the outer solar atmosphere. How rapidly does the spatial distribution of the Poynting flux vary in time?  In contrast to the snapshot we analyze here, successive Poynting flux maps would be needed to address this question.  While SP vector magnetograms are the best currently available, the telemetry limitations of {\\em Hinode} preclude long-duration runs of successive, rapid rasters over moderately large FOVs.  Consequently, the HMI instrument aboard SDO \\citep{Scherrer2012} could be used investigate the temporal variation of the Poynting flux. Unfortunately, HMI has both worse spatial resolution and poorer spectral sampling.  So a related question is: How sensitively do estimates of the Poynting flux depend upon a magnetograph's spatial and spectral resolution?  Analysis of a region simultaneously observed with SOT and HMI would be worthwhile. (It is probable, in fact, that additional energy flux could be resolved with even higher-resolution observations, though the energy flux must begin to decrease at some limit, to avoid an ultraviolet catastrophe.  This motivates studies with new, higher-resolution instruments, as discussed below.)  Within the larger context of the coronal heating problem, we suggest that a key strategic observational objective for understanding chromospheric and coronal heating should be construction of a detailed energy budget for the photosphere-to-corona system, with spatially and temporally resolved energy inputs correlated with energy release in all forms --- radiation, kinetic energy in thermal and non-thermal particles and bulk motion, and gravitational potential energy.  This will require high-resolution and high-cadence observations of the magnetic field and emission throughout the photosphere-to-corona system, for which both space-based observatories (e.g., SDO, IRIS, and the planned Solar-C satellite\\footnote{http://hinode.nao.ac.jp/SOLAR-C/Documents/Solar-C\\_e.pdf}) and existing and planned ground-based observatories (NST [Goode \\etal 2010], GREGOR [Volkmer \\etal \\, 2010], ATST [Rimmele \\etal \\, 2010], and EST [Zuccarello and Zuccarello 2011]) will be essential. \\nocite{Goode2010} \\nocite{Volkmer2010} \\nocite{Rimmele2010} \\nocite{Zuccarello2011}   Acknowledgments: We thank the {\\em Hinode} science teams for their hard work in producing the excellent data that made this study possible. We thank: the referee for constructive criticisms that we believe improved the manuscript; Anthony Yeates for pursuing estimation of the Poynting flux in unipolar regions, which inspired this work; NWRA's K. D. Leka for help with SP data prepared by the late Tom Metcalf; Bruce Lites for providing error estimates from the ASP inversions of the SP data and suggestions regarding interpretation of the data; and T.D. Tarbell and G.H. Fisher for providing useful comments about the manuscript that helped improve it.  B.T.W. also gratefully thanks the Japan Society for the Promotion of Science, whose fellowship supported much of the work that underlies this study, and acknowledges funding from the NSF's National Space Weather Program under award AGS-1024862, the NASA Living-With-a-Star TR\\&T Program (grant NNX11AQ56G), and the NASA Heliophysics Theory Program (grant NNX11AJ65G).  The authors are grateful to Japanese and U.S.  taxpayers for providing the funds necessary to perform this work.  {\\em Hinode} is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, and NASA and STFC (UK) as international partners. Scientific operation of the {\\em Hinode} mission is conducted by the {\\em Hinode} science team organized at ISAS/JAXA. This team mainly consists of scientists from institutes in the partner countries. Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (UK), NASA (USA), ESA, and NSC (Norway)."
    },
    "1402/1402.2792_arXiv.txt": {
        "abstract": "We  study  orbital inclination  changes  associated with the precession  of a disc-planet system that occurs  through  gravitational interaction with a  binary companion  on an inclined orbit. We  investigate  whether this  scenario  can account for   giant planets on close  orbits highly inclined  to the stellar equatorial plane.  We   obtain  conditions for  maintaining  approximate coplanarity and test them  with   SPH-simulations.   For  parameters of interest, the  system undergoes  approximate rigid body precession  with    modest warping while the  planets migrate inwards.  Because of  pressure forces, disc self-gravity is not needed   to maintain the configuration.  We consider a disc and single  planet for different  initial  inclinations of the binary orbit to the midplane of the   combined system and  a system of three planets for which  migration leads to  dynamical instability  that reorders the planets. As the interaction  is dominated by  the time averaged quadrupole  component of the binary's  perturbing potential, results for  a circular  orbit can be scaled to apply to  eccentric orbits. The  system  responded adiabatically  when  changes to binary  orbital parameters occurred  on  time scales exceeding the  orbital period. Accordingly  inclination changes are  maintained under its slow removal. Thus the  scenario for generating high inclination  planetary orbits   studied here, is promising. ",
        "introduction": "In the past few decades, the number of detected extrasolar planets around other main sequence stars has increased dramatically, so that  at the time of writing,  more than 1000 planets have been discovered. Well studied  planet formation  scenarios  such as  core accretion  \\citep{Miz1980, Pol1996} or disc fragmentation \\citep{May2002} involve the planet forming in a disc with the natural expectation that the orbit should be  coplanar. However, Rossiter McLaughlin measurements  of close orbiting hot Jupiters indicate that around $40\\%$ have angular momentum vectors  significantly misaligned with  the angular velocity vector of the central star \\citep[e.g.][]{Tri2010, Alb2012}. As the stellar and disc angular velocities are naturally expected to be aligned, this would imply an inclination of the planet orbit relative to the nascent protoplanetary disc.   This has been underlined by several recent studies of the interactions of planets with a range of initial orbital inclinations with respect to a disc \\citep[e.g.][]{Cre2007, Mar2009, Bit2011}. In these works the evolution of planets with masses  up to a maximum of  $1\\rmn{M_J}$ with different initial eccentricities and relatively small initial inclinations up to a maximum of $15^\\circ$ interacting with three-dimensional isothermal and radiative discs are considered. They showed damping for both eccentricity and inclination with the planets circularising  in the disc after a few hundred orbits. Planets with higher inclinations have also been studied by e.g. \\cite{ Rei2012}  and  \\cite{Tey2013}. A detailed survey of a range of planet masses and the full range of orbital inclinations interacting with a disc  has been presented by \\cite{Xia2013} (hereafter Paper 1). The timescale of realignment with the gas disc was found to be  comparable with the disc  lifetime for very high inclinations $ >60^{\\circ}$ and planet masses of  one Jupiter mass. For smaller initial relative inclinations  and/or  larger planet masses,  the timescale for  realignment is shorter than the disc lifetime. These results taken together imply that there is  a need for an explanation for the origin  of  misaligned giant planets, if it is assumed that the discs angular momentum vector is always  aligned with the spin axis of the central star.      Several  scenarios have been proposed to explain the origin of misalignments between the planetary orbital angular  momentum vector  and the stellar rotation  axis  as well as produce close in giant planets. The first involves excitation  of very high eccentricities, either through  the Lidov-Kozai effect induced by the  interaction with a distant companion \\citep[e.g.][]{Fab2007, Wu2007}, or through planet-planet scattering or chaotic interactions \\citep[e.g.][]{Weid1996,Rasio1996,Pap2001,Nag2008}. When the planet attains a small enough pericentre distance, this is then followed by orbital circularization due to tidal interaction with the central star, leaving the planet on a close inclined  circular orbit.  A method for producing high disc inclinations with respect to the stellar equator has also been indicated by \\cite{Tho2003}, who  have studied the evolution of two giant planets in resonance, adopting an approximate analytical expression for the influence of a gas disc in producing  orbital migration. Their calculations suggest that resonant inclination excitation can occur when  the eccentricity of the inner planet reaches a threshold value  $\\sim  0.6$. Inclinations gained by the resonant pair of planets can reach values up to $\\sim 60^\\circ$.  However in  a recent study,   \\citet{Daw2013} find that large eccentricities,  that are expected to  be associated with such interactions for  orbits too wide to be affected by stellar tides,   are seen  mostly in metal rich systems. Accordingly they conclude  that  gentle  disc migration and planet-planet scattering must  both operate during the early evolution of giant planets. Only systems packed with giant planets, which most easily form around metal rich stars, can produce large eccentricities through scattering. They also note a lack of a correlation between spin-orbit misalignment and metallicity that  may indicate the operation of a mechanism that can cause misalignment of the disc while simultaneously  allowing disc migration of giant planets into close orbits.    An explanation for the origin of  misaligned planets, that is consistent  with formation in  and inward migration driven by a disc may be connected with the possibility that the orientation of the disc changes,  possibly due  either to  a lack of orbital  alignment of the source material or to gravitational encounters with passing stars. This may occur   either before or after  planet formation  \\citep[e.g.][]{Bat2010,Thi2011}.  \\cite{Thi2011}  showed that the capture of gas from accretion envelopes by  a protoplanetary disc could cause it to become significantly misaligned with respect to its original plane. Although no planets are involved in their simulations, it is of   interest to remark that  misaligned gas discs can easily be produced when  interactions  with the environment out of which they form are considered. In this regard we note that  protostars are not generally isolated but are  bound in stellar clusters where gravitational interactions  are common.   Recently, the system Kepler-56 has been studied by \\cite{Hub2013} who used asteroseismological methods to determine the obliquity in that system. Kepler-56 has  two known transiting planets. \\cite{Hub2013} found that the planets have angular momentum vectors  strongly inclined  with  respect to the central stellar spin axis. Taking into account the large observed obliquities,  dynamical simulations with the two known planets  failed to produce a stable non coplanar  planetary system. The authors then studied an alternative scenario involving a distant companion and were able to show in dynamical simulations that they succeeded in producing the architecture of the Kepler-56 system including  coplanarity of the planetary orbits  and the high obliquity. Further observations are required to specify the mass of the companion which is crucial to determine  the origin of the misalignment in Kepler-56. In the case of a planetary companion, the misalignment is inferred to occur after the planets have  formed while in the case of a stellar companion, the misalignment is inferred to  occur while  the host gaseous disc is present.   In this  context,  \\cite{Bat2012} suggested that a binary companion  could be responsible for  producing  misalignments between stellar spin axes and planetary orbital planes as a result of making  the protoplanetary disc precess. The process envisaged the formation and disruption of wide binaries  in birth clusters. that dissolve on timescales up to a few tens of millions of years. The excitation of significant misalignment should  then occur within the same time scale.  We remark that the interaction between a binary companion and a disc could cause the disc to become warped. Notably, \\cite{Ter2013} studied the effect of a warped disc on the orbit of a Jupiter mass  planet. For a sufficiently warped disc, it was shown that  a  high orbital inclination  for the  planet  might be excited through direct gravitational interaction.   In view of the above discussion, it is important to investigate the response of a system composed of both a protoplanetary disc and  one or more planets that can induce a gap and migrate within it,  under the gravitational perturbation of a binary companion in a misaligned orbit. If  misalignment  with the stellar equatorial plane is induced during the late stages,  planets may reside in a cavity in the inner regions of the disc, relatively little further accretion of disc material onto the central star may take place optimising the final misalignment,  thus making this an interesting case to study.  In this paper we consider the response of the  disc-planet system as a whole obtaining  conditions,  based on estimates of precessional torques,  for it to maintain approximate coplanarity and test these against the results of numerical simulations.  From these  we are able to show that for appropriate parameters of intererest, the whole system precesses approximately as a rigid body with the disc exhibiting only a modest amount of warping. This does not   necessitate a self-gravitating disc as assumed by \\citet{Bat2012} because the disc can communicate  more effectively through the action of pressure forces \\citep[eg.][]{Lar1996}  than through density waves as long as the Toomre stability parameter exceeds unity,  a situation  expected  for the most part during the later stages of protoplanetary discs   \\citep[see][for a detailed discussion]{Papl99}. Even if the outer parts of the disc are maintained close to marginal stability, effects due to self-gravity and pressure are  expected to be comparable for global warps. Accordingly we make the simplification of neglecting the disc self-gravity which results in a very significant reduction in computational resource requirements. We also test how the  response adjusts to changes  in the orbital parameters of the binary finding that this is essentially adiabatic even when the changes occur on a timescale comparable to the binary orbital period.  It has been shown that Smoothed particle hydrodynamics (SPH) simulations are capable   of being applied to the problem of planet-disc interaction \\citep[e.g.][]{Sch2004,Val2006} though in comparison to grid-based simulation methods, they are generally more diffusive and incur greater computational expense. But in contrast to grid-based methods, as they adopt a Lagrangian approach, SPH simulations can be readily used to simulate a gas disc  with a free boundary   that undergoes a large amount of movement  in three dimensions with the disc having the freedom to change its shape at will. As it is suitable for a system precessing with high mutual inclinations, we adopt this approach  here.   In Section \\ref{sec:sim_details}, we   describe our simulation technique, giving details  of the modified locally isothermal equation of state we adopted in Section \\ref{sec:eos} and details  of the smoothing length and artificial viscosity prescriptions in Section \\ref{sec:artv}. In Section \\ref{sec:IC}, we describe the general setup for  the disc,  the  planets, the central star and the binary companion. In Section \\ref{sec:theory}, we give a brief theoretical overview of the  global interaction between the planet-disc system and a   binary star with orbital plane  misaligned with the  disc midplane. Estimates of the precessional torques that can act  between different components in the system are given   in   appendix \\ref{ap:precession}. If the putative precession frequencies they could produce, significantly exceed the putative precession frequency of  the entire system calculated assuming that it rotates as a rigid body, these torques can be utilised to maintain alignment. Precessional  Torques acting between a disc and a planet in a gap/inner cavity as well as between two planets without  a disc are considered.  We go on to present  the results of our simulations   in  Section \\ref{sec:1pl}. Simulations with a disc together with a single migrating giant planet and a binary companion in an inclined orbit  are discussed in Section \\ref{sec:1p2}. These  systems are found to maintain approximate alignment consistently with  estimated magnitudes of internal precessional  torques. The adiabatic response to changes in the binary orbit is discussed in Section \\ref{Adi}. A simulation with a disc, misaligned binary orbit and a three planet system is described in Section \\ref{sec:3pl}. The disc-planet system  is seen to maintain approximate coplanarity, as a result of disc-planet torques, as  in the single planet case, even though there is a phase of dynamical instability. Finally, we sumarize and discuss our results in Section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}   We have performed SPH-simulations in order to study the influence of a  stellar binary companion on a circumstellar disc and massive planets that may orbit  therein. The gas disc and the planets were initiated on coplanar  circular orbits around the central star.  We studied  four inclinations  of the plane of the binary orbit  to the initial disc midplane.  These were    $i_B=[30^{\\circ}, 45^{\\circ},60^{\\circ} {\\rm and}\\,  80^{\\circ}]$. The masses of the central star and binary companion were taken to be $1\\ \\rmn{M}_{\\odot}.$ We considered planets of mass $2\\ \\rmn{M_J}$ and $1\\ \\rmn{M_J}$  in a disc of mass $0.01\\ \\rmn{M_\\odot}$ with outer radial boundary at $20$~AU. The binary orbit was taken to be circular with orbital radius $100$~AU.  The central finding was that the planet and the gas disc evolved  maintaining  approximate coplanarity for the full range of $i_B$. Apart from the geometric configuration of the orbits, this coplanarity was manifested in diverse ways,  such as the invariance of the inward planet  migration or the gap profile. Importantly and quite generally, the relative inclination of the disc and planet in respect to the binary star was constant through the simulations. The enforcement of coplanarity means that large inclination changes between the plane of the disc-planet system and the original disc plane, which might coincide with  the stellar equatorial  plane, can attain large values, especially for large $i_B.$  Such changes could be made permanent if the binary companion is removed.  We remark that we have neglected any  precessional or accretion torques acting on the central star which might change the orientation of  its equatorial plane.  However, \\citet{Bat2012}  showed that precessional torques could be neglected for a slowly rotating star when the disc extended down to  the corotation radius.  Furthermore, they are expected to be significantly weaker for a disc with an inner cavity as envisaged here. Accretion torques are also likely to be weak in these circumstances and notably \\citet{Bat2013} indicate that these depend on the form of the stellar magnetic field and do not need to act in a way that would tend to make the system align.  Although the specific choice of simulation parameters relating to the initial disc size and binary orbit were chosen largely for numerical convenience, we point out that the results can be extended to apply to other configurations in a number of ways.  Apart from  being able to apply the usual scaling of lengths and times that leave gravitationally interacting systems invariant, we remark that the  important  perturbation component  from the binary companion that is responsible for  precession and orbital orientation changes is the time averaged quadrupole term that is dominant  for large orbital separation. This term takes the same form regardless of the eccentricity of the orbit, provided $D$ is replaced by $A_B(1-e_B^2),$ where $A_B$ and $e_B$ are the semi-major axis and eccentricity of the binary orbit respectively. Thus, as long as the quadrupole term is dominant we expect similar results. In addition, we estimated the precession period from our  simulations  to be  $3.8\\times 10^4(20{\\rm AU}/R_{out})^{3/2}(D/100{\\rm AU})^{3/2}\\ \\rmn{yr}$ (see Section \\ref{precto}). Thus the same precession period is obtained if $R_{out}$ and $D$ are increased by the same factor.  Although the disc is larger, the dynamics in the inner region should not be affected as long as internal communication in the disc allows quasi-rigid body precession. Note too that if we increase $D$ without increasing $R_{out},$ the precession period increases making the maintenance of coplanarity even easier.    In summary, the main factors affecting the maintenance of coplanarity are the ability of the disc to communicate within a precession period and the ability of the components to exert strong enough precessional torques. These can be estimated from equations (\\ref{Torque4})-(\\ref{w2pp}) and will be seen to be determined by the disc and planet parameters within a radial scale defined by the planetary system rather than the outer boundary radius. This generates many possibilities for producing significant orbital reorientation through interaction with the binary.   Furthermore  we also ran test simulations in order to model the slow removal of the binary companion. By doing this, we noted the adiabatic behaviour of the system. We found a smooth transition of the system to  one with only a central star. Due to the lack of the external torque, the system stopped its precession,  attaining a  constant inclination between the midplane of the disc-planet  system  and the original disc midplane. This indicates that large orbital orientation changes could be produced through a sequence of interactions with different binary companions as long as the dynamical processes associated with evolution of the binary orbits is long compared to the characteristic dynamical time.    We also studied a multi-planet system with three planets and disc under the influence of a binary companion. We performed a comparison between the N-body/SPH code GADGET-2 and  a 2D grid based code  NIRVANA  for the system with no binary companion. The planets with masses 0.4, 1 and 2 $\\rmn{M_J}$ were put on coplanar circular orbits near to a 2:3:6 resonance. Even though there are significant  differences in implementation  such as the treatment of the boundary conditions  and  the application and behaviour  of viscosity,  the form of the migration  of the planets is in  satisfactory agreement.  When   the  planetary system  was embedded in a gas disc in the presence of a binary companion, it was found that inward migration resulted in the appearance of instability. After 500 orbits, the innermost planet   of  0.4 $\\rmn{M_J}$  was scattered to  attain the outermost orbit after which the evolution stabilized  with the planets again attaining  near circular orbits. Throughout this evolution,  the  orbital  inclinations  of all three planets remained strongly coupled to the  disc inclination. In this respect the multi-planet system behaved exactly as the single planet system.   Thus the simulations that we have performed have shown that when the disc is induced to undergo quasi-rigid body precession, as long as this is on a long enough timescale that internal precessional torques can retain coplanarity. This is maintained when there is a migrating planetary system with gaps in the interior part of the disc.   In this regard,  we remark that \\cite{Ghe1993} have conducted an observational survey of T Tauri stars in two neighbouring star forming regions and obtained a binary  frequency in the  separation range 16 - 252 AU for T Tauri stars of $\\sim$ 60 \\% which is roughly 4 times greater than the binary star frequency for solar-type main-sequence stars. As  this frequency  can be altered  by evolutionary processes within  star clusters which can form  and disrupt binary systems,  \\cite{Ghe1993} proposed that the high binary star frequency associated with  T Tauri stars is correlated with their young ages  and is  thus an evolutionary effect. However, this  could be crucial  for  a  protoplanetary disc, since as shown in our simulations, an inclined binary companion is able to produce a disc with a high inclination with respect to the stellar equatorial plane in a time  that is short compared to the disc's lifetime. More recent studies of binaries in young  stellar systems mainly support the results of \\cite{Ghe1993}  \\citep[e.g.][]{ Koe2001, Kra2007,  Kro2011}. Although there is a large uncertainty, observations reveal a frequency of binary systems in stellar clusters of 30 $\\sim$ 50 \\% \\citep{Duc2004, Kra2009, Mar2012}. Note that this is comparable to the frequency of hot Jupiters with orbits that are highly inclined with respect to the stellar equatorial plane. Given the high  abundance of binary stellar systems and the short timescale required for inclination generation, a scenario involving  a binary companion with separation  of the order of 100 AU   could accordingly  give  a possible explanation for the observed frequency of planets on highly  inclined orbits."
    },
    "1402/1402.2101_arXiv.txt": {
        "abstract": "We use Principal Component Analysis (PCA) to study the gas dynamics in numerical simulations of typical MCs. Our simulations account for the non-isothermal nature of the gas and include a simplified treatment of the time-dependent gas chemistry. We model the CO line emission in a post-processing step using a 3D radiative transfer code. We consider mean number densities $n_0 = 30, 100, 300\\,$cm$^{-3}$ that span the range of values typical for MCs in the solar neighbourhood and investigate the slope $\\alpha_\\text{PCA}$ of the pseudo structure function computed by PCA for several components: the total density, H$_2$ density, $^{12}$CO density, $^{12}$CO $J = 1 \\rightarrow 0$ intensity and $^{13}$CO $J = 1 \\rightarrow 0$ intensity. We estimate power-law indices $\\alpha_\\text{PCA}$ for different chemical species that range from 0.5 to 0.9, in good agreement with observations, and demonstrate that optical depth effects can influence the PCA. We show that when the PCA succeeds, the combination of chemical inhomogeneity and radiative transfer effects can influence the observed PCA slopes by as much as $\\approx \\pm 0.1$. The method can fail if the CO distribution is very intermittent, e.g. in low-density clouds where CO is confined to small fragments.  \\endabstract  \\begin{keywords} galaxies: ISM -- ISM: clouds -- ISM: molecules \\end{keywords} ",
        "introduction": "\\label{sec:introduction}  Assessing the dynamics of molecular clouds (MCs) is crucial for understanding star formation. Besides magnetic fields, gravity, chemistry and the thermodynamics of MCs, turbulence plays a fundamental role in the star formation process \\citep[see e.g.,][]{MacLowAndKlessen2004,ElmegreenAndScalo2004,McKeeAndOstriker2007,HennebelleAndFalgarone2012}. Turbulent shocks can create (or disperse) local overdense regions or filaments within MCs that may contract to pre-stellar cores and finally form stars.  The characterization of turbulence in the ISM is necessary to understand the star formation process, the structure of the ISM itself and the distribution of turbulent kinetic energy over different spatial scales. In an extremely influential paper, \\citet{Larson1981} noted that the then-available observational data suggested that the velocity dispersion of the gas within molecular clouds was correlated with the size of the cloud. He found that the CO linewidth $\\delta v$ scaled with the size $\\ell$ of a molecular cloud, as $\\delta v \\propto \\ell^{0.38}$. Later studies \\citep[e.g.][]{SolomonEtAl1987} found a slightly steeper relationship, $\\delta v \\propto \\ell^{0.5}$. Furthermore, several studies revealed that MCs in other galaxies \\citep[e.g.][]{BolattoEtAl2008} and in the ISM in the Galactic center \\citep[e.g.][]{OkaEtAl2001,ShettyEtAl2012} also follow a similar relationship. During the last decades, this linewdith-size relationship has been the focus of several studies \\citep[e.g.][]{OssenkopfAndMacLow2002,HeyerAndBrunt2004,HeyerEtAl2009,RomanDuvalEtAl2011}.  There are several established methods for measuring velocity fluctuations as a function of spatial scale, e.g. by measuring the spectral correlation function \\citep{RosolowskyEtAl1999}, the velocity channel analysis \\citep{LazarianEtAl2000,LazarianEtAl2004} or the $\\Delta$-variance method \\citep{OssenkopfEtAl2008a,OssenkopfEtAl2008b}. Another well-established method is Principal Component Analysis (PCA), applied in this context by \\citet{HeyerAndSchloerb1997,BruntAndHeyer2002a}. PCA is a multivariate mathematical tool that reduces redundant information of a dataset by approximating a large number of statistical variables with a smaller number of significant linear combinations. In this way, PCA can provide a description of turbulent motion by measuring velocity fluctuations as a function of spatial scale. Linear regression then yields a slope of the pseudo structure function obtained from observational measurements \\citep[see e.g.,][]{HeyerAndSchloerb1997,BruntAndHeyer2002b,HeyerAndBrunt2004,HeyerEtAl2006,RomanDuvalEtAl2011} or from numerical simulations or analysis of synthetic images \\citep[see e.g.,][]{BruntAndHeyer2002a,BruntEtAl2003,FederrathEtAl2010}.  Theoretical efforts are necessary to understand the influence of different effects on the PCA slopes, such as geometric projection \\citep[see e.g.,][]{OstrikerEtAl2001,Ballesteros-ParedesAndMacLow2002,ShettyEtAl2010}, the radiation field, optical depth effects \\citep{OssenkopfAndKlessenAndHeitsch2001}, density fluctuations \\citep{RomanDuvalEtAl2011} or the non-uniform chemical makeup of the gas. Chemical effects may be important for the determination of velocity fluctuations as a function of spatial scale, since the relation between H$_2$ and CO is not linear. Furthermore, CO accurately traces the total gas only in dense regions, and so CO emission may not accurately reveal the properties of the diffuse gas (\\citealt{GloverEtAl2010}; \\citeauthor{ShettyEtAl2011a}~2011a).  Recent developments have allowed us to use high-resolution, three-dimensional chemistry models in numerical simulations of turbulent interstellar gas \\citep{GloverEtAl2010} to get a better insight into how the choice of chemical species as a tracer of gas affects the slope of the pseudo structure function within MCs. Time-dependent chemistry is important for modelling the formation and destruction of molecules in diffuse environments. Radiative transfer calculations are then required in order to convert the simulation output into synthetic emission maps. We focus in this paper on how chemistry and radiative transfer affect the observed slopes of the pseudo structure function in different environments, spanning a range of possible densities representative of MCs. We also compare the theoretical findings with recent observational efforts by \\citet{RomanDuvalEtAl2011}.  In section \\ref{sec:methodsandsims} we describe the simulations, the radiative transfer post-processing and the PCA technique, which we use to compute the pseudo structure functions. In section \\ref{sec:results} we present the results of our PCA of simulations with different mean number densities. We discuss our findings in section \\ref{sec:discussion}, and present our conclusions in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We applied PCA to numerical simulations of MCs with time dependent chemistry to describe the impact of chemistry, radiative transfer and of the mean density on the turbulent dynamics of MCs. We studied the PPV cubes of five components: the total number density, H$_2$, $^{12}$CO (each without radiative transfer) and $^{12}$CO and $^{13}$CO, respectively (each with radiative transfer). In each case, we analysed the PCA pseudo structure function for different initial number densities $n_0$. We computed individual slopes $\\alpha_\\text{PCA}$ and compared them to observational measurements. Furthermore, we also studied the influence of radiative transfer on the results of the PCA. We report the following findings:  \\begin{enumerate} \\item $^{12}$CO is a good tracer of H$_2$ in compact regions in MCs, but not in low density regions. H$_2$ is more extended than the $^{12}$CO density as it is better able to self-shield in diffuse regions. \\item Optical depth effects can be important for the PCA structure analysis. Since $^{12}$CO is optically thick and can be easily excited, there is a much higher degree of coherence in the spatial dimension of MCs for the $^{12}$CO intensity than for the $^{12}$CO density, as shown in Section \\ref{opticaldepth}. Large scales are therefore suppressed in the $^{12}$CO density. Since $^{13}$CO is generally optically thin, it traces the structure of a molecular clouds more accurately. \\item We are able to reproduce observations by performing simulations with chemistry and radiative transfer postprocessing for $^{12}$CO and $^{13}$CO, as shown in Figure \\ref{fig:hist}. We measure slopes $\\alpha_\\text{PCA}$ of the pseudo structure function in a range from 0.5 to 0.9. The slope distributions in Figure \\ref{fig:hist} peak at $0.59 \\pm 0.02$ for $^{12}$CO intensity for the n100 and at $0.82 \\pm 0.03$ and $0.74 \\pm 0.02$ for the n300 model for $^{12}$CO and $^{13}$CO intensity, respectively. \\item We compute mean values of $\\langle l_1 / l_0 \\rangle \\gtrsim 0.2$ that are consistent with previous studies of \\citet{OssenkopfAndMacLow2002} and \\citet{BruntEtAl2009} and show that the driving scale is not significantly affected by the chemistry and the radiative transfer. \\item We analyse the log-density dispersion $\\sigma_s$ of the different PDFs, finding $\\sigma_s < 2$ for the total density and H$_2$ density field. For the CO density field we find $\\sigma_s > 2$ which is due to the high degree of chemical inhomogeneity. \\end{enumerate}  Despite the difference in the morphology evident in the maps (Figure \\ref{fig:images} and \\ref{fig:spectra}), most of the resulting PCA slopes for all components are only marginally different, depending on density. For the higher density model n300, where the effects from stellar feedback that are not modeled and are likely to affect the results, the turbulence statistics of $^{12}$CO and $^{13}$CO intensity provide reasonably good approximations to the underlying structure of the density and velocity field. For the fiducial n100 model, which has a lower density, differences between the underlying physical parameters and the emergent CO intensity lead to more notable discprencies in the corresponding $\\alpha_\\text{PCA}$.  For further investigations and to get a better understanding of how different initial conditions (e.g.\\ different metallicities and chemical abundances, turbulent driving on different scales and amplitudes, different radiation fields, etc.) affect the slopes $\\alpha_\\text{PCA}$, we need numerical simulations that span a wider range of physical parameters. Varying the turbulence driving on different spatial scales could be necessary to model local feedback from other stars or from stellar winds."
    },
    "1402/1402.1933_arXiv.txt": {
        "abstract": "We present a derivation of the cosmological distance-redshift relation up to second order in perturbation theory. In addition, we find the observed redshift and the lensing magnification to second order. We do not require that the density contrast is small, we only that the metric potentials and peculiar velocities are small. Thus our results apply into the nonlinear regime, and can be used for most dark energy models. We present the results in a form which can be readily computed in an N-body simulation. This paper accompanies Paper I, where the key results are summarised in a physically transparent form and applications are discussed. ",
        "introduction": "The next generation of cosmological surveys will map the universe to extremely high precision across a vast range of scales. These surveys may be sensitive to a variety of subtle nonlinear relativistic effects~-- most of which remain to be analyzed. The angular diameter-redshift relation, which is a key determinant of any cosmological model, is affected by inhomogeneities along the line of sight. Nonlinear contributions change apparent sizes and brightnesses of objects, beyond the usual convergence contributions from standard gravitational lensing~\\cite{Bartelmann:1999yn} and from the relativistic linear correction in the form of Doppler lensing~\\cite{Bonvin:2008ni,Bacon:2014uja}. These corrections could well be significant: for example, a deep spherical void can change the magnitude-redshift relation by up to 20\\% beyond the prediction of linear perturbation theory when all relativistic effects are taken into account~\\cite{Bolejko:2012uj}. Accurately quantifying the amplitude of the corrections around realistic structures is important for the precise modelling now required for large-scale surveys.  Light  propagation  in inhomogeneous  spacetimes  gives  rise to modification of the area distance  due to both Ricci focussing (from the matter along the ray) and Weyl focussing (from the tidal effects of nearby matter). It leads to corrections in redshift due to the differences between the true (observed) redshift of a source and its redshift in an average smoothed-out model. For light propagating in a perturbed Friedmann-Lemaitre-Robertson-Walker (FLRW) space-time, linear perturbations of the angular diameter, or area, distance were computed first by Sasaki \\cite{Sasaki:1987}, in synchronous gauge, and then by Pyne and Birkinshaw \\cite{Pyne:2003bn} in the Poisson gauge.  Bonvin, Durrer and Gasparini \\cite{Bonvin:2005ps} carefully included all relativistic effects for scalar perturbations at first order, and this was extended to include vector and tensor modes in~\\cite{DiDio:2012bu}.    Partial results towards the full second-order computation  of cosmological distance were given by Barausse, Matarrese and Riotto~\\cite{Barausse:2005nf}, in the case of a matter-dominated universe. The full second-order result, for a universe with dark energy, was presented independently in our paper I \\cite{Umeh:2012pn} and by Ben-Dayan et al~\\cite{BenDayan:2012wi} (see also related work~\\cite{BenDayan:2012ct, BenDayan:2013gc,Fanizza:2013doa,Nugier:2013tca}).  Our result, when specialized to the partial formulas presented in \\cite{Barausse:2005nf} and transformed to the comoving-synchronous gauge used in \\cite{Barausse:2005nf}, is in agreement with \\cite{Clarkson:2011uk}. The task of establishing agreement between our version of the full results and the version of \\cite{BenDayan:2012wi} is highly nontrivial, because their work was done not only in another gauge, but also in a different coordinate system (geodesic lightcone coordinates). Here we give the detailed derivation of the results presented in paper I.        We assume a flat FLRW background; the extension to curved backgrounds is nontrivial and is left for future work. Vector and tensor modes at first order are neglected. However, we do incorporate second-order vector and tensor modes that are sourced by the product of first-order scalar fluctuations ~\\cite{Lu:2007cj,Ananda:2006af,Lu:2008ju,Baumann:2007zm}. For the vector modes, this is particularly important because the amplitude of the vector potential can be of order 1\\% of the first-order scalar potential. Tensor modes are included at second order for completeness. We allow for dynamical dark energy, with only one restriction, i.e. that we neglect anisotropic stress at first order. This excludes only a small subset of dark energy models. (Note that at first order after decoupling, anisotropic stress in CDM and baryons is negligible, while for photons and neutrinos the anisotropic stress makes a negligible contribution.) Anisotropic stress is generated at second order, and we include this.  In general relativity, the field equations at first order imply equality of the metric perturbations in Poisson gauge, $\\Psi=\\Phi$, when the anisotropic stress vanishes. We adopt $\\Psi=\\Phi$ as an assumption at first order -- but we do not make any further use of the field equations at first or second order. Our perturbative expansions assume that the metric perturbations and peculiar velocity of matter are small, but we do not require the density contrast to be small, so that our results apply into the mildly nonlinear regime.  The assumption that $\\Psi=\\Phi$ at first order means that our results do not apply to general modified gravity theories that are an alternative to dark energy. When $\\Psi\\neq\\Phi$ at first order, there are significant complications, and we leave this more general case for future work.  There are a variety of ways to calculate our result. We work mainly in a perturbed Minkowski spacetime, conformally transforming our key results at the end to FLRW spacetime. The key steps in our derivation are as follows (perturbatively evaluated order by order): \\begin{enumerate} \\item Solve the null geodesic equation for the photon 4-momentum as a function of background affine parameter. \\item Calculate the redshift  as a function of affine parameter. \\item Solve the Sachs equations for the area distance as a function of affine parameter. \\item[$\\rightarrow$] All quantities are now in terms of the affine parameter of the background. The distance-redshift relation at this stage is in the form of two parametric equations with the background affine parameter linking them. \\item Perturbatively invert the redshift-affine parameter relation and substitute into the area-distance relation. \\end{enumerate} We then find an explicit expression for the area distance to a source at observed redshift $z_s$ as $D_A(z_s)$, given by~\\eqref{bigboy}, which is written as a function of the background comoving distance $\\chi_s$ to the source -- which is in turn determined using the background distance-redshift relation for the \\emph{observed} redshift (more on this confusing issue in Section~\\ref{sec:obs}). We write our result in terms of the metric potentials and peculiar velocity, all defined in the Poisson gauge.  This paper is organised as follows: we provide general  covariant evolution equations for null shear, null expansion and the area distance in Section \\ref{sec:Covframe-work}. The photon geodesic equations are derived and solved perturbatively in Section \\ref{sec:geodesic-pert},  and we calculate the physical redshift in terms of the perturbed metric variables in Section \\ref{sec:physical-red}. The area distance in perturbation theory is derived and solved order by order in Section \\ref{sec:area-ditance}. We present a general nonlinear expression for expressing the area distance in terms of the physical redshift of the source in Section \\ref{sec:obs}, which contains our main result (an alternative form is to be found in  Appendix~\\ref{alternative}).  {\\em Notation:} We use indices $a,b,c,\\cdots=0,1,2,3$ in a general spacetime. In perturbed FLRW, the indices $i,j,\\cdots=1,2,3$ denote spatial components, and the linearly perturbed metric in Poisson gauge (scalar modes) is $a^2[-(1+2\\Phi)\\d\\eta^2+(1-2\\Phi)\\delta_{ij}\\d x^i\\d x^j]$. The conformal Hubble rate is $\\mathcal{H}=a'/a$. ",
        "conclusions": "We have presented the full expression for the observed area distance-redshift relation to second-order on an arbitrary flat FLRW background~\\eqref{bigboy}. Other than our assumption that $\\Phi=\\Psi$ at first-order, our formula is valid for any dark energy model~-- excluding only dark energy models with anisotropic stress. % We have not assumed explicitly that over-densities must be small. Consequently, our result should hold into the mildly nonlinear regime.  We have presented the result in a ready-to-use form. Given a first-order potential $\\Phi$, and a matter model, the second-order potential and velocities are fixed through the Einstein field equations~-- see e.g., \\eqref{phi2}-- \\eqref{delta2} in the case of LCDM. Each term in~\\eqref{bigboy} is then directly computable. This can be implemented in N-body codes to correct for line-of-sight effects. Alternatively, it can be used to compute the spherical harmonic expansion of $\\langle\\hat D_A(\\hat z_s)\\rangle$, which then gives the deviation from the background geometry expected in the standard model. This latter `backreaction' effect has been postulated to be both large and small, and is very important to calculate accurately, in order to correctly determine the background geometry and parameters (see e.g., \\cite{Rasanen:2011bm,Clarkson:2011zq,Clarkson:2011uk,Marozzi:2012ib,BenDayan:2012ct}).     \\subsection*{ Acknowledgments}   We thank Camille Bonvin and Ruth Durrer for useful discussions and  Giovanni Marozzi and Gabriele Veneziano for comments. Most of the computations here were done with the help of the tensor algebra packages xAct /xPert~\\cite{Brizuela:2008ra} and  xPand~\\cite{Pitrou:2013hga}. OU and RM are supported by the South African Square Kilometre Array (SKA) Project, CC and RM are supported by the National Research Foundation (South Africa), RM is supported by the Science \\& Technology Facilities Council (UK) (grant no. ST/K00090X/1), and all authors were supported by a Royal Society (UK)/ NRF (SA) exchange grant.  \\appendix"
    },
    "1402/1402.1758_arXiv.txt": {
        "abstract": "We derive the equations of celestial mechanics governing the variations of the orbital elements under a stochastic perturbation generalizing the classical Gauss equations. Explicit formulas are given for the semi-major axis, the eccentricity, the inclination, the longitude of the ascending node, the pericenter angle and the mean anomaly which are express in term of the angular momentum vector $\\textbf{H}$ per unit of mass and the energy $E$ per unit of mass. Together, these formulas are called the \\emph{stochastic Gauss equations} and they are illustrated numerically on an example from satellite dynamics. ",
        "introduction": "Nowadays celestial mechanics is used by a wide class of scientists which provide multiple applications (see \\cite{murray_dermott}, \\cite{burns1976} and references therein). In all these works, the underling nature of the model considered is always deterministic. However, considering models with randomness or stochastic behavior is not an easy problem (examples for celestial mechanics can be found in \\cite{cr}, \\cite{cpp} and \\cite{bcp}). Indeed, the nature and the origin of such a model needs a real discussion of the phenomena that we want to study.\\\\  Most of the problems in celestial mechanics are seen as a two-body problem perturbed by a force. For example, the main approach of the $n$-body problem is to consider two bodies, in mutual gravitational interaction, which are perturbed by the other bodies. In that case, the perturbed force is the gravitational attraction of the other bodies.\\\\  When we are dealing with more than two bodies, or more generally with an arbitrary perturbing force, the orbital elements, which characterize the trajectory of the bodies, do not remain constant. In that case, the main tool of celestial mechanics to study the perturbed problem, is the set of equations given the variations of the orbital elements called, the \\textit{Gauss equations}. Because the Gauss equations allow studying general problems in celestial mechanics, we propose in this paper to generalize them to the stochastic case which include by definition the deterministic case. \\\\  We follow the strategy of \\cite{burns1976} who derived Gauss's equations for the elliptical case with elementary considerations which defined the orbital elements in mean of the angular momentum per unit of mass and the energy per unit of mass. From an example of the satellite dynamics, we illustrate numerically the variation of the orbital elements associated. Finally, we give the variation of the Laplace-Runge-Lenz vector. It allows deriving the variation of the orbital elements in more general cases. For example, in the cases of null inclination, hyperbolic or parabolic configurations. ",
        "conclusions": "In this article, we have developed the stochastic perturbation equations of celestial mechanics which generalize the classical Gauss equations. This is done with the It\\^o theory of stochastic differential equations and with basic considerations on the angular momentum and the energy per unit of mass. This approach allows predicting the impact of each components of the stochastic perturbing force on the dynamic. From a perturbing acceleration containing white noises, we showed the construction of the stochastic perturbation associated and we illustrated numerically the dynamic associated with the stochastic Gauss equations. Finally, we derived the variation of the Laplace-Runge-Lenz vector in order to obtain the minimum set of equations covering a large class of problem in celestial mechanics for further studies and applications."
    },
    "1402/1402.3794_arXiv.txt": {
        "abstract": "{The NASA \\kepler\\ mission has been continuously monitoring the same field of the sky since the successful launch in March 2009, providing high-quality stellar lightcurves that are excellent data for asteroseismology, far superior to any other observations available at the present. In order to make a meaningful analysis and interpretation of the asteroseismic data, accurate fundamental parameters for the observed stars are needed. The currently available parameters are quite uncertain as illustrated by e.g. \\cite{thygesen}, who found deviations as extreme as 2.0 dex in \\feh\\ and \\logg, compared to catalogue values. Thus, additional follow-up observations for these targets are needed in order to put firm limits on the parameter space investigated by the asteroseismic modellers. Here, we propose a metod for deriving accurate metallicities of main sequence and subgiant solar-like oscillators from medium resolution spectra with a moderate S/N. The method takes advantage of the additional constraints on the fundamental parameters, available from asteroseismology and multi-color photometry. The approach enables us to reduce the analysis overhead significantly when doing spectral synthesis, which in turn will increases the efficiency of follow-up observations.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.6452_arXiv.txt": {
        "abstract": "Standard-model extension (SME) is an effective field theory introducing all possible Lorentz-violating (LV) operators to the standard model (SM) and general relativity (GR). In the pure-gravity sector of minimal SME (mSME), nine coefficients describe dominant observable deviations from GR. We systematically implemented twenty-seven tests from thirteen pulsar systems to tightly constrain eight linear combinations of these coefficients with extensive Monte Carlo simulations. It constitutes the first detailed and systematic test of the pure-gravity sector of mSME with the state-of-the-art pulsar observations. No deviation from GR was detected. The limits of LV coefficients are expressed in the canonical Sun-centered celestial-equatorial frame for convenience of further studies. They are all improved by significant factors of tens to hundreds with existing ones. As a consequence, Einstein's equivalence principle is verified substantially further by pulsar experiments in terms of local Lorentz invariance in gravity. ",
        "introduction": "The transformation matrix ${\\cal R}$ from the Sun-centered celestial-equatorial coordinate system, $(\\hat{\\bf X},\\hat{\\bf Y},\\hat{\\bf Z})$, to the binary coordinate system, $(\\hat{\\bf a},\\hat{\\bf b},\\hat{\\bf c})$, can be decomposed into five simple parts, \\begin{equation} \\left( \\begin{array}{c} \\hat{\\bf a} \\\\ \\hat{\\bf b} \\\\ \\hat{\\bf c} \\end{array} \\right) = {\\cal R} \\left( \\begin{array}{c} \\hat{\\bf X} \\\\ \\hat{\\bf Y} \\\\ \\hat{\\bf Z} \\end{array} \\right) \\,, \\end{equation} with \\begin{equation} {\\cal R} = {\\cal R}^{(\\omega)} {\\cal R}^{(i)} {\\cal R}^{(\\Omega)} {\\cal R}^{(\\delta)} {\\cal R}^{(\\alpha)} \\,. \\end{equation} For a binary pulsar, explicit expressions of these matrices are \\begin{eqnarray} {\\cal R}^{(\\alpha)} &=& \\left( \\begin{array}{ccc} -\\sin\\alpha & \\cos\\alpha & 0 \\\\ -\\cos\\alpha & -\\sin\\alpha & 0 \\\\ 0 & 0 & 1 \\end{array} \\right) \\,,\\\\ {\\cal R}^{(\\delta)} &=& \\left( \\begin{array}{ccc} 1 & 0 & 0 \\\\ 0 & \\sin\\delta & \\cos\\delta \\\\ 0 & -\\cos\\delta & \\sin\\delta \\end{array} \\right) \\,,\\\\ {\\cal R}^{(\\Omega)} &=& \\left( \\begin{array}{ccc} \\cos\\Omega & \\sin\\Omega & 0 \\\\ -\\sin\\Omega & \\cos\\Omega & 0 \\\\ 0 & 0 & 1 \\end{array} \\right) \\,,\\\\ {\\cal R}^{(i)} &=& \\left( \\begin{array}{ccc} 1 & 0 & 0 \\\\ 0 & \\cos i & \\sin i \\\\ 0 & -\\sin i & \\cos i \\end{array} \\right) \\,,\\\\ {\\cal R}^{(\\omega)} &=& \\left( \\begin{array}{ccc} \\cos\\omega & \\sin\\omega & 0 \\\\ -\\sin\\omega & \\cos\\omega & 0 \\\\ 0 & 0 & 1 \\end{array} \\right) \\,, \\end{eqnarray} where $\\alpha$ (right ascension), $\\delta$ (declination), $\\Omega$ (longitude of ascending node), $i$ (orbital inclination), $\\omega$ (longitude of periastron) are defined in Figs.~1--2 in the main text.  For a solitary pulsar, only a rotation from $(\\hat{\\bf X},\\hat{\\bf Y},\\hat{\\bf Z})$ to $(\\hat{\\bf I},\\hat{\\bf J},\\hat{\\bf K})$ is needed, hence one can set ${\\cal R}^{(\\Omega)}$, ${\\cal R}^{(i)}$, and ${\\cal R}^{(\\omega)}$ to unit matrices. ",
        "conclusions": "}  In the parametrized post-Newtonian (PPN) formalism, there are two parameters, $\\alpha_1$ and $\\alpha_2$, together with an absolute velocity with respect to a preferred frame, ${\\bf w}$, to describe semiconservative LV effects~\\cite{will93,will06}.  We notice that if we do the replacements in the pure-gravity sector of mSME, $\\bar s^{0j} \\to \\alpha_1 w^j/4$ and $\\bar s^{jk} \\to - \\alpha_2 w^j w^k$, the dynamics of PPN is exactly recovered. We can clearly see the relation from Lagrangians or from a comparison between Eqs.~(\\ref{eq:edot}--\\ref{eq:ldot}) and Eqs.~(13--19) in~\\cite{sw12}. This new relation was also discussed by Bailey and Kosteleck\\'y in a different manner, see Eq.~(68) in~\\cite{bk06}.  In scalar-tensor theories, strong fields associated with neutron stars may develop nonperturbative effects~\\cite{de93}. One would expect similar effects in SME. Therefore our results in Table~\\ref{tab:LV} are, strictly speaking, {\\it effective} constraints that include strong-field contributions.  The tests here are of probabilistic nature because of unknown angles $\\Omega$ and $\\eta$.  High-precision timing of nearby pulsars (e.g. PSR~J0437$-$4715~\\cite{vbs+08}) and observations of interstellar scintillation (e.g. PSR~B1534+12~\\cite{bplw02}) may determine $\\Omega$. When more such measurements are available, robust tests can be performed.  All tests in this Letter are based on time derivatives of observables, whose precisions improve as $T_{\\rm obs}^{-3/2}$. Continuous observations can quickly improve all tests.  Especially in the era of FAST and SKA telescopes, more pulsars and high sensitivities will further enable better tests of gravity.  We also notice that some timing results used here were reported more than one decade ago. New results for sure have already improved a lot. We hope observers report $\\dot e$ and $\\dot x$ measurements whenever feasible, which can give more realistic constraints for LV coefficients. It is also possible to use dissipative effects~\\cite{ybby13} once gravitational radiation is calculated in SME.  Recently, Bailey et al.~\\cite{beo13} obtained the first constraint of the $\\bar s^{\\rm TT} = \\bar s^{\\rm XX} + \\bar s^{\\rm YY} +\\bar s^{\\rm ZZ}$ component from Gravity Probe B. Their result $|\\bar s^{\\rm TT}| < 3.8 \\times 10^{-3}$ (68\\% CL) can be combined with ours to break all degeneracy. We also note that the boost between a pulsar system and the Solar System can be used to mix $\\bar s^{\\rm TT}$ with other components; it hence provides the possibility to obtain a constraint of $\\bar s^{\\rm TT}$ directly from pulsars. Currently, only a few pulsars with optical observations of their companions have 3\\,D velocity measurements that are needed for the boost calculation. We expect more measurements will achieve the goal in the future."
    },
    "1402/1402.3041_arXiv.txt": {
        "abstract": "The XMM-Newton spectral-fit database is an ongoing ESA funded project aimed to construct a catalogue of spectral-fitting results for all the sources within the XMM-Newton serendipitous source catalogue for which spectral data products have been pipeline-extracted ($\\gtrsim$ 120,000 X-ray source detections). The fundamental goal of this project is to provide the astronomical community with a tool to construct large and representative samples of X-ray sources by allowing source selection according to spectral properties. ",
        "introduction": "The XMM-Newton serendipitous source catalogue is the largest catalogue of X-ray sources built to date, see \\cite[Watson \\etal\\ (2009)]{watson}. In its latest version, the 3XMM-Newton Data Release 4 (3XMM DR4\\footnote{http://xmmssc-www.star.le.ac.uk/Catalogue/3XMM-DR4/}), it contains photometric information for more than 500,000 source detections, corresponding to $\\sim$ 370,000 unique sources. Besides, spectra and time series were also extracted if the number of source counts collected by the EPIC camera was $>$ 100 counts. The resulting catalogue contains spectra for more than 120,000 detections corresponding to $\\sim$ 85,000 unique sources. The project described here, is aimed to take advantage of the great wealth of data and information contained within the XMM-Newton source catalogue, to construct a database composed of spectral-fitting results. ",
        "conclusions": ""
    },
    "1402/1402.3055_arXiv.txt": {
        "abstract": "One of the most spectacular predictions of the general theory of relativity is the black hole, an object that plays a central role in modern physics [1,2,3] and astrophysics [4,5] Black holes are, however, plagued by fundamental paradoxes that remain unresolved to this day. First, the black hole event horizon is teleological in nature [6] which means that we need to know the entire future space-time of the universe to determine the current location of the horizon. This is essentially impossible.  Second, any information carried by infalling matter is lost once the material falls through the event horizon. Even though the black hole may later evaporate by emitting Hawking radiation [7], the lost information does not reappear, which has the rather serious and disturbing consequence that quantum unitarity is violated [8]. Here we propose that the above paradoxes are restricted to a particular idealized model of collapse first studied in the 1930s [9, 10] in which the event horizon, which defines the boundary of the black hole, forms initially, and the singularity in the interior of the black hole forms at a later time. In contrast, gravitational collapse from more reasonable and/or physically more realistic initial conditions often leads to models in which the event horizon and the singularity form simultaneously.  We show that this apparently simple modification mitigates the causality and teleological paradoxes and at the same time lends support to two recently proposed solutions to the information paradox, namely, the ``firewall'' [11] and ``classical chaos'' [12]. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.4991_arXiv.txt": {
        "abstract": "We present a {\\it Revised IRAS-FSC Redshift Catalogue} (RIFSCz) of 60,303 galaxies selected at 60\\ $\\micron$ from the {\\it IRAS Faint Source Catalogue} (FSC). This revision merges in data from the WISE All-Sky Data Release, the Tenth SDSS Data Release (DR10), the GALEX  All-Sky Survey Source Catalog (GASC), the 2MASS Redshift Survey (2MRS) and the Planck Catalogue of Compact Sources (PCCS). The RIFSCz consists of accurate position, ultra-violet (UV), optical, near-, mid- and far-infrared, sub-millimetre (sub-mm) and/or radio identifications, spectroscopic redshift (if available) or photometric redshift (if possible), predicted far-infrared and sub-mm fluxes ranging from 12 to 1380\\ $\\micron$ based upon the best-fit infrared template. We also provide stellar masses, star-formation rates and dust masses derived from the optical and infrared template fits, where possible. $56\\%$ of the galaxies in the RIFSCz have spectroscopic redshifts and a further $26\\%$ have photometric redshifts obtained through the template-fitting method. At S60 $>$ 0.36 Jy, the 90$\\%$ completeness limit of the FSC, $93\\%$ of the sources in the RIFSCz have either spectroscopic or photometric redshifts. An interesting subset of the catalogue is the sources detected by Planck at sub-mm wavelengths. 1200 sources have a detection at better than 5-$\\sigma$ in at least one Planck band and a further 1186 sources have detections at 3-5$\\sigma$ in at least one Planck band. % ",
        "introduction": "The {\\it IRAS Faint Source Catalog} (FSC; Moshir et al. 1992) contains 173,044 sources reaching a depth of $\\sim$0.2 Jy at 12, 25 and 60\\ $\\micron$. It is limited to $|b|>20^\\circ$ in unconfused regions at 60\\ $\\micron$. For sources with high-quality flux density\\footnote{In the IRAS FSC, the flux density quality (FQUAL) is classified as high (=3), moderate (=2) or upper limit (=1)}, reliability is $>99\\%$ at 12 and 25\\ $\\micron$ and $>94\\%$ at 60\\ $\\micron$. Around 41$\\%$ of the FSC sources are detected at 60 $\\mu$m (FQUAL $>$ 1). The construction of  the Imperial IRAS-FSC Redshift Catalogue (IIFSCz; Wang \\& Rowan-Robinson 2009) was made possible by overlaps (in terms of depth and area) with various imaging and/or redshift 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 6dF Galaxy Survey (Jones et al. 2004, 2005). However, the large positional error of the IRAS sources has meant in particular that the cross-identification between IRAS and the deep optical data SDSS is very challenging. The completion and/or release of several major surveys in the last few years, including the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) All-Sky Survey, the GALEX All-Sky Survey Source Catalog (GASC; Seibert et al., in prep.), the 2MASS Redshift Survey (2MRS; Huchra et al. 2012), the Tenth SDSS Data Release (DR10; Ahn et al. 2013) and the Planck all-sky survey (Planck Collaboration I 2013) in the microwave and sub-millimetre (sub-mm), has made it imperative to revise the IIFSCz. In particular, the imaging depth and sky coverage of WISE are such that the WISE all-sky source catalogue provides counterparts for the majority of sources in the IIFSCz and thereby improves the positional accuracy of these sources by a factor of $\\sim$10. Using the WISE positions, cross-identification between IRAS sources and sources detected at UV, optical and near-infrared becomes much easier and much more reliable. %    The layout of this paper is as follows. In Section~\\ref{sec:identification}, first we cross-identify IRAS FSC sources with their mid-infrared (MIR) WISE counterparts using a likelihood ratio technique. Then, using their WISE positions, we cross-match FSC sources with sources detected in other wavebands including the ultra-violet (UV), optical, near-infrared (NIR) and sub-mm.  Spectroscopic redshifts are collected from a number of databases, e.g., SDSS DR10, 2MRS and NED. In Section 3, we estimate photometric redshift using a template-fitting method for sources matched with WISE, SDSS and/or 2MASS counterparts but do not have spectroscopic redshifts. In Section 4,  infrared templates are fitted to the mid- and far-infrared data from IRAS and WISE as well as sub-mm data from Planck for FSC sources with WISE and Planck associations. Once a best-fit infrared template is found, we make flux predictions at far-infrared and sub-mm wavelengths.  Finally,  discussions and conclusions of the overall properties of the revised IRAS-FSC Redshift  catalogue (RIFSCz) are given in Section 5. Throughout the paper, we adopt a flat cosmological model with $\\Lambda=0.7$ and $h_0=0.72$ and a Salpeter initial mass function (Salpeter et al. 1955). Unless otherwise stated, we use the AB magnitude system, and $\\log=\\log_{10}$. ",
        "conclusions": "To summarise, the RIFSCz contains a total of 60,303 galaxies, $56\\%$ of which have spectroscopic redshifts from NED, FSSz, PSCz, 6dF and the SDSS spectroscopic DR10 survey and $26\\%$ of which have photometric redshifts from the template-fitting method. At a flux limit of S60 = 0.36 Jy which is the $90\\%$ completeness limit of the FSC, more than $93\\%$ of the galaxies in the RIFSCz have either spectroscopic or photometric redshifts.The RIFSCz covers about $61\\%$ of the whole sky. In Fig.~\\ref{skydistribution}, the sky distribution in Galactic coordinates is plotted for all galaxies in the RIFSCz with either spectroscopic or photometric redshift and galaxies without any redshift estimate.  The RIFSCz catalogue represents the best all-sky far-infrared (60\\ $\\micron$) selected galaxy catalogue with multi-wavelength photometric information. As a result, several extragalactic science programmes can be carried out with the RIFSCz dataset. For example, we can examine the relation between the ratio of the infrared to UV luminosity  (the infrared excess: IRX) and the slope of the UV spectrum ($\\beta$). This relation is critical to estimating the dust attenuation in galaxies and hence the true star-formation rate. In the top panel in Fig.~\\ref{irx_beta}, we plot IRX versus $\\beta$ for RIFSCz galaxies with GALEX UV detections at SNR $\\ge5$ and the reduced $\\chi^2$ of the infrared SED fit $\\le3$. The typical statistical error of the infrared luminosity ($L_{\\rm IR}$) for these galaxies is about 0.3 dex.  In the bottom panel in Fig.~\\ref{irx_beta}, we plot IRX versus $\\beta$ for a subset of the galaxies in the top panel which are detected at 350\\ $\\micron$ by Planck. The typical statistical error of $L_{\\rm IR}$ for this subset of galaxies with 350\\ $\\micron$ is around 0.1 dex. Our $\\beta$ values are derived using the recipe from Kong et al. (2004). A compilation of the IRX - $\\beta$ relations published in the literature is also plotted. The Calzetti et al. (2000) and Overzier et al. (2011) relations are derived from UV-selected samples. The Boquien et al. (2012) relation is derived from a $K$-band selected sample. The Mu{\\~n}oz-Mateos et al. (2009) relation is derived from the SINGS sample (Kennicutt et al. 2003) which consists of 75 local galaxies spanning a large range in morphology, luminosity, SFR and opacity.  The Hao et al. (2011) relation is derived from the SINGS sample combined with a sample of local galaxies covering the full range of optical spectral characteristics (Moustakas \\& Kennicutt 2006). It is clear that there is a large dispersion among the different relations derived from samples with different selection criteria and our far-infrared selected sample show a general good agreement with these relations. In addition, our far-infrared selected sample also seems to cover regions with high IRX but low $\\beta$ values which are not represented in other relations.  Detailed investigations are required to fully understand the causes of the difference between our IRX - $\\beta$ relation based on far-infrared selection and the other relations. However, it is beyond the scope of this paper.   We provide a long version and a short version of the RIFSCz to the public, available from \\url{http://astro.ic.ac.uk/public/mrr/fss/readme} and \\url{http://www.astro.dur.ac.uk/~lwang83/RIFSCz.tar.gz}. The long version contains all the information (e.g. positional, photometric or spectroscopic) we have either assembled or derived in the process of constructing the RIFSCz. Then short version contains the most important information of the sources in the RIFSCz and should suffice for most users. The format of the long and short version of the RIFSCz catalogue is explained in detail in the Appendix.   We have presented the revised IRAS-FSC Redshift Catalogue (RIFSCz). It contains 60,303 galaxies selected at 60\\ $\\micron$ from the IRAS Faint Source Catalog, covering around $61\\%$ of the whole sky. The process of retrieving spectroscopic redshifts, multi-wavelength source identification and photometric redshift estimation is described in detail.  Users should be aware of issues such as the intrinsic IRAS FSC completeness limit, the redshift completeness variations across the sky and the varying quality of photometric redshifts derived for different subsets of the catalogue.    \\begin{figure*} \\includegraphics[height=2.95in,width=5.4in]{sky_distribution_specz.pdf} \\includegraphics[height=2.95in,width=5.4in]{sky_distribution_photz.pdf} \\includegraphics[height=2.95in,width=5.4in]{sky_distribution_noz.pdf} \\caption{The sky distribution (in Galactic coordinates) of all galaxies in the RIFSCz which have received spectroscopic redshifts (top), all galaxies with photometric redshift but without spectroscopic redshifts (middle), and all galaxies with neither spectroscopic or photometric redshifts (bottom).} \\label{skydistribution} \\end{figure*}   \\begin{figure} \\includegraphics[height=3.in,width=3.45in]{irx_beta_subset1.pdf} \\includegraphics[height=3.in,width=3.45in]{irx_beta_subset2.pdf} \\caption{The infrared to UV luminosity (the infrared excess: IRX) versus the slope of the UV spectrum ($\\beta$) for galaxies in the RIFSCz catalogue (top: galaxies with SNR $\\ge5$ UV detections and the reduced $\\chi^2$ of the infrared SED fit $\\le3$; bottom: a subset of the galaxies in the top panel which are detected at 350\\ $\\micron$.), compared with other relations published in the literature. The typical statistical error of the infrared luminosity ($L_{\\rm IR}$) for galaxies in the top and bottom panel is about 0.3 dex and 0.1 dex respectively.} \\label{irx_beta} \\end{figure}"
    },
    "1402/1402.6736_arXiv.txt": {
        "abstract": "We analyze SDSS spectra of 568 obscured luminous quasars. The [OIII]$\\lambda$5007\\AA\\ emission line shows blueshifts and blue excess, indicating that some of the narrow-line gas is undergoing an organized outflow. The velocity width containing 90\\% of line power ranges from 370 to 4780 km/sec, suggesting outflow velocities up to $\\sim$2000 km/sec, and is strongly correlated with the radio luminosity among the radio-quiet quasars. We propose that radio emission in radio-quiet quasars is due to relativistic particles accelerated in the shocks within the quasar-driven outflows; star formation in quasar hosts is insufficient to explain the observed radio emission. The median radio luminosity of the sample of $\\nu L_{\\nu}$[1.4GHz]$=10^{40}$ erg/sec suggests a median kinetic luminosity of the quasar-driven wind of $L_{\\rm wind}=3\\times 10^{44}$ erg/sec, or about 4\\% of the estimated median bolometric luminosity $L_{\\rm bol}=8\\times 10^{45}$ erg/sec. Furthermore, the velocity width of [OIII] is positively correlated with mid-infrared luminosity, which suggests that outflows are ultimately driven by the radiative output of the quasar. Emission lines characteristic of shocks in quasi-neutral medium increase with the velocity of the outflow, which we take as evidence of quasar-driven winds propagating into the interstellar medium of the host galaxy. Quasar feedback appears to operate above the threshold luminosity of $L_{\\rm bol}\\sim 3\\times 10^{45}$ erg/sec. ",
        "introduction": "\\label{sec:intro}  Black hole feedback -- the strong interaction between the energy output of supermassive black holes and their surrounding environments -- is routinely invoked to explain the absence of overly luminous galaxies, the black hole vs. bulge correlations and the similarity of black hole accretion and star formation histories \\citep{tabo93, silk98, spri05, hopk06}. After years of intense observational effort, specific examples of black-hole-driven winds have now been identified using a variety of observational techniques, both at low and at high redshifts \\citep{nesv06, arav08, nesv08, moe09, dunn10, alex10, harr12, fabi12}.  How these outflows are launched near the black hole and established over the entire host galaxy remains a topic of active research. In particular, it is becoming clear that as the outflow impacts an inhomogeneous interstellar medium of the galaxy, the winds are expected to contain gas at a wide range of physical conditions, and these different ``phases'' of the winds require different types of observations \\citep{veil05}. As a result, the determination of the physical parameters of these outflows -- including such basic parameters as the mass, the momentum and the kinetic energies they carry -- remains challenging. In the last few years, a lot of progress in this area has been made by observations of the coldest components of the outflows which are in the form of neutral or even molecular gas \\citep{fisc10, feru10, rupk11, stur11, aalt12, veil13, rupk13a, rupk13b, cico14, sun14}. It remains unclear how common this cold component is in winds driven by a powerful active nucleus and how the mass, momentum and energy carried by the wind are distributed across the different phases.  Several years ago, we embarked on an observational program to determine whether radio-quiet, luminous quasars have observable effects on their galaxy-wide environment. One of our lines of investigation is to determine the extent and kinematics of the warm ($T\\sim 10^4$ K) ionized gas -- the so-called narrow-line region of quasars. We simplify the observational task by looking at obscured quasars \\citep{zaka03, reye08} -- those where the line of sight to the nucleus is blocked by intervening material, allowing us to study the distribution of matter in the galaxy unimpeded by the bright central source. We find strong evidence that ionized gas is extended over scales comparable to or exceeding that of the host galaxy; furthermore, it is kinematically disturbed and is not in equilibrium with the gravitational potential of the galaxy \\citep{gree09, gree11, hain13}.  More recently, we surveyed a sample of obscured radio-quiet quasars using a spectroscopic integral field unit \\citep{liu13a, liu13b}. We found extended ionized gas encompassing the entire host galaxy (median diameter of nebulae of 28 kpc), suggestive of wide-angle outflows, and determined kinetic energies of these outflows to be well in excess of $10^{44}$ erg/sec, with a median 2\\% conversion rate from the bolometric luminosity to the kinetic energy of warm ionized gas; more energy can be carried by other components. Furthermore, we identified several candidate objects where the wind has ``broken out'' of the denser regions of the galaxy and is now expanding into the intergalactic medium, sometimes in bubble-like structures \\citep{gree12, liu13b}.  These observations demonstrate the presence of extended ionized gas in host galaxies of type 2 quasars, which is apparently out of dynamical equilibrium with the host galaxy and is likely in an outflow on the way out of the host. In this paper, we examine spectra of several hundred obscured quasars and we study the relationships between gas kinematics and other physical properties of these objects. In Section \\ref{sec:data} we describe the sample selection, the dataset and the measurements. In Section \\ref{sec:optical}, we conduct kinematic analysis of the optical emission lines. In Section \\ref{sec:multiwv} we discuss the relationships between multi-wavelength properties of quasars and  kinematic measures of their ionized gas nebulae. In Section \\ref{sec:compo} we present composite spectra and discuss trends in weak emission lines. We present qualitative models for radio emission and emission lines in Section \\ref{sec:discussion}, and we summarize in Section \\ref{sec:conclusions}.  We use a $h$=0.7, $\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7 cosmology throughout this paper. SDSS uses vacuum wavelengths, but for consistency with previous literature we use air wavelengths in Angstroms to designate emission lines. Wavelengths are obtained from NIST \\citep{kram13} and Atomic Line List\\footnote{\\url{http://www.pa.uky.edu/~peter/atomic/}} and converted between air and vacuum as necessary using \\citet{mort91}. Objects are identified as SDSS Jhhmm+ddmm, with full coordinates given in the catalog by \\citet{reye08}. We use `1D', `2D' and `3D' abbreviations for one-, two-, and three-dimensional values. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{Gas kinematics and outflows}  In this paper we study the kinematics of ionized gas emission in 568 luminous obscured quasars from the sample of \\citet{reye08}, primarily using their [OIII]$\\lambda\\lambda$4959,5007 emission lines. For every object, we determine a set of non-parametric measures of line asymmetry, velocity width and shape. We find that objects with blueshifted emission and line asymmetries are more prevalent in our sample, which we take as a signature that outlows in which the redshifted part is affected by dust extinction are common in our sample. The velocity widths we see in our sample (median $w_{80}=752$ km/sec and $w_{90}=1060$ km/sec, max $w_{80}=2918$ km/sec and $w_{90}=4780$ km/sec) are much higher than those in starburst galaxies (median $w_{90}\\simeq 600$ km/sec, \\citealt{rupk13a,hill14}).  Using the conversion between the line-of-sight velocity width of the line and the outflow velocity from Section \\ref{sec:outflows}, we can estimate outflow velocities as $v_0\\simeq w_{80}/1.5$. We find that for about half of the objects in our sample, $v_0$ ranges from $\\sim$500 to $\\sim$2000 km/sec. It is likely that in every object a range of cloud velocities is present; these estimates should be thought of as the median velocities in each source. From the observations presented here, it is not known which spatial scales dominate these estimates, but from our spatially resolved observations \\citep{liu13b} it is clear that at least some of the clouds maintain similar velocities all the way out to $\\sim 10$ kpc from the quasar.  Neither the overall width nor the width of the narrower component in the multi-Gaussian line decomposition correlate with the velocity dispersion of the host galaxy. We therefore find no evidence for significant amounts of gas in dynamical equilibrium with the host (e.g., a component associated with the rotating galaxy disk). The higher velocity gas tends to be on average blueshifted by $\\sim 100$ km/sec relative to the lower velocity gas. This may mean that the higher velocity gas is found on somewhat smaller spatial scales which are more prone to dust extinction. This conclusion is in line with the slight observed decrease of line-of-sight velocity dispersions of the gas as a function of the distance from the nucleus seen in spatially resolved observations of quasar winds \\citep{liu13b}. The effect is very subtle (3\\% decrease per each projected kpc), and overall high gas velocities are maintained over the entire host galaxy \\citep{liu13b}. Other strong lines, such as [OII] and H$\\beta$, also show outflow signatures, but significantly more mild than those seen in the [OIII] line. In the objects with the most extreme [OIII] kinematics, [OII] is much narrower than [OIII], with the kinematics of H$\\beta$ being in between the two.  The increase in shock-diagnostic lines [SII]$\\lambda\\lambda$4069,4076 and [NI]$\\lambda\\lambda$5198,5200 as a function of [OIII] width strong support for using the [OIII] width as a proxy for the outflow velocity. [OII] declines both with [OIII] width and [OIII] luminosity, suggesting that the [OII]-emitting regions are over-ionized by photo-ionization and quenched by increasing shock velocity. We find evidence for shock-ionization contribution to coronal lines [FeVII].  \\subsection{Outflows and radio emission}  [OIII] velocity width correlates strongly with radio luminosity. We suggest that the radio emission is a by-product of the outflow activity, with particles accelerated on the shock fronts as the radiatively driven quasar wind propagates into the interstellar medium of the host galaxy (this mechanism was first hypothesized by \\citealt{stoc92} to explain the low levels of radio emission in broad absorption line quasars). The median radio luminosity in our sample, $\\nu L_{\\nu}$[1.4GHz]$=10^{40}$ erg/sec, requires kinetic energy of the outflow of $L_{\\rm wind}=3\\times 10^{44}$ erg/sec if the efficiency of conversion in quasar-driven winds is similar to that in supernova-driven winds. The bolometric luminosities of our objects are poorly known, but we estimate that the median value for our sample is $8\\times 10^{45}$ erg/sec and thus the kinetic luminosity of the wind is 4\\% of the bolometric power. This idea is further reinforced by the approximately quadratic dependence $\\nu L_{\\nu}$[1.4GHz]$\\propto$(width[OIII])$^2$, which implies a linear relationship between radio luminosity and kinetic energy of the outflow. [OIII] velocity is positively correlated with the infrared luminosity, which suggests that the ionized gas outflow is ultimately driven by the radiative pressure close to the accreting black hole.  Another possible explanation for the [OIII] width / radio correlation is that the mechanical energy of the relativistic jet (which has not yet broken out of its host galaxy) is used to heat an overpressured cocoon \\citep{bege89} which then launches the wind of ionized gas \\citep{mull13}. As long as the jet is still compact, it appears as an unresolved radio core in FIRST observations. We use the scalings between the core radio luminosity and the jet kinetic energy by \\citet{merl07} to estimate how much kinetic power would be required to produce the observed median radio luminosity. We find $L_{\\rm jet}=3\\times 10^{44}$ erg/sec, a value which is almost identical to the wind kinetic power obtained in the previous paragraph, even though they were estimated using completely different methods. This is likely not a coincidence: the efficiency of conversion of mechanical luminosity into synchrotron radiation is determined by the fraction of energy that can be converted into relativistic particles on shock fronts, which is likely independent of the origin of these shocks.  The similarity of energy requirements underscores the difficulty of distinguishing between these two mechanisms. In the jet scenario one expects to see a jet in the high-resolution radio observations, whereas in the wind scenario radio emission is more diffuse and is present everywhere where shocks are propagating; however, in practice the collimated part of the jet does not dissipate energy very efficiently and so can be hard to detect. In both models, the morphology of the ionized gas depends much more strongly on the distribution of the interstellar medium than on the exact driving mechanism \\citep{gaib11, wagn12, wagn13}. Thus, the morphology of ionized gas is not necessarily a useful clue.  Radio spectral index could be a useful measurement for identifying recent or on-going particle acceleration (flatter synchrotron spectra for freshly accelerated, more energetic particles), but again particle acceleration may be happening in both scenarios. Interestingly, \\citet{lal10} find flatter spectral indices in type 2 quasars than expected for jets in the standard geometry-based unification models, so a combination of spectral and morphological investigations in the radio may be worth pursuing further.  One uncomfortable consequence of the jet scenario is that it requires for every radio-quiet quasar to have a powerful jet, with only a minority of them being active long enough to to break out of the galaxy (otherwise there would be too many extended radio sources). Another problem is that the jet scenario provides no ready explanation for the correlation between [OIII] width and infrared luminosity, unless jets are also directly connected to the accretion power \\citep{falk95}.  In the wind scenario, if the initial distribution of matter is roughly spherically symmetric, so will be the ionized gas emission and the radio emission. In a disk galaxy, once the size of the wind reaches the vertical scaleheight of the disk it propagates largely along the path of least resistance perpendicular to the disk \\citep{wagn13}, forming two symmetric bubbles on either side. Ionized gas emission is concentrated in shells and filaments, whereas the radio emission is filling the bubbles. In low-resolution data, one would see both the line emission and the radio emission oriented roughly in the same direction, so distinguishing this mechanism from jet-induced outflow requires high-quality observations.  Such bubbles are directly seen in the radio in external galaxies \\citep{ceci01, hota06} and in our own Milky Way, where the bubbles are known as ``the microwave haze'' \\citep{fink04,su10}. These structures are also seen in X-rays which often closely trace the morphology of the radio emission and / or the ionized gas filaments \\citep{ceci01, cros08, wang09}. In our Galaxy not only X-rays are observed \\citep{sofu00}, but the structures as seen in gamma-rays and are known as ``the Fermi bubbles'' \\citep{su10}. For comparison, the mechanical luminosity necessary to inflate the bubbles in NGC 3079 \\citep{ceci01} is $\\sim 30$ times lower than our median $L_{\\rm wind}$, whereas that inferred for ``the Fermi bubbles'' in the Milky Way is $10^4-10^5$ times lower than our $L_{\\rm wind}$. It is interesting that all of the examples above host jets \\citep{ceci01, cros08, wang09, su12}, but it remains unclear whether this is universally true and whether they are contributing most of the required power.  \\subsection{Feedback in low- and high-luminosity active galaxies}  Our investigation focuses on the most luminous type 2 quasars at $z\\la 1$. In these sources, we find strong evidence for quasar-driven winds on galaxy-wide scales, for radio emission associated with these winds, and for ionized gas velocities in excess of escape velocity from the galaxy. In lower luminosity active galaxies, many previous studies demonstrated that ionized gas is in dynamical equilibrium with the host galaxy and that radio emission is likely a bi-product of star-formation processes. Thus quasar-driven feedback may be present above some threshold in luminosity and absent below this threshold. We estimate this threshold by observing that in Figure \\ref{pic_wise}, our sources separate well from star-forming galaxies at about $\\nu L_{\\nu}$[12\\micron]$=2\\times 10^{44}$ erg/sec. Using bolometric corrections from \\citet{rich06}, \\citet{liu09} and \\citet{liu13b}, we estimate the corresponding bolometric luminosity to be $L_{\\rm bol}=3\\times 10^{45}$ erg/sec (uncertainty $\\pm 0.4$ dex).  Given the crudeness of this estimate, we consider it to be consistent with the value of $L_{\\rm bol}=2.4\\times 10^{45}$ erg/sec (uncertainty $\\pm 0.3$ dex) suggested by \\citet{veil13} based on incidence of molecular outflows in ultraluminous infrared galaxies \\citep{veil13}. In principle the threshold value should be dependent on the depth of the potential (and therefore the stellar velocity dispersion), as well as the amount of gas that needs to be accelerated \\citep{zubo12}, and it will be interesting to see whether these ideas are borne out in future analyses."
    },
    "1402/1402.1320_arXiv.txt": {
        "abstract": "{ In the framework of the unification scheme of radio-loud active galactic nuclei, BL Lac objects and quasars are the beamed end-on counterparts of low-power (FR\\,I) and high-power (FR\\,II) radio galaxies, respectively. } { Some BL Lacs have been found to possess the FR\\,II-type large-scale radio morphology, suggesting that the parent population of BL Lacs is a mixture of low- and high-power radio galaxies. This seems to apply only to `low frequency-peaked'  BL Lacs, since all the `high frequency-peaked' BL Lacs studied so far were shown to host exclusively the FR\\,I-type radio jets. While analyzing the {\\emph {NVSS}} survey maps of the TeV detected BL Lacs, we have however discovered that the high frequency-peaked object SHBL\\,J001355.9$-$185406 is associated uniquely with the one-sided, arcmin-scale, and edge-brightened jet/lobe-like feature extending to the south-west from the blazar core. } { In order to investigate in detail the large-scale morphology of SHBL\\,J001355.9$-$185406, we have performed low-frequency and high-resolution observations of the source at 156, 259 and 629\\,MHz using the Giant Metrewave Radio Telescope. } { Our analysis indicates that no diffuse arcmin-scale emission is present around the unresolved blazar core, and that the arcmin-scale structure seen on the {\\emph {NVSS}} map breaks into three distinct features unrelated to the blazar, but instead associated with background AGN. } { The upper limits for the extended radio halo around the TeV-emitting BL Lac object SHBL\\,J001355.9$-$185406 read as $< 10\\%-1\\%$ at $156-629$\\,MHz. The fact that the integrated radio spectrum of the unresolved blazar core is flat down to 156\\,MHz indicates that a self-similar character of the jet in the source holds up to relatively large distances from the jet base. } ",
        "introduction": "\\label{intro}  Blazars constitute the most extreme class of active galactic nuclei (AGN), for which the total radiative output is dominated by a Doppler-boosted and highly variable emission of a nuclear relativistic jet observed at a small viewing angle. The blazar family includes flat-spectrum radio-loud quasars (FSRQs) and BL Lacertae (BL Lac) objects. In the framework of the ``unified scheme'' of AGN, powerful radio galaxies of the Fanaroff-Riley morphological type II (FR\\,IIs; see, \\citealt{1974MNRAS.167P..31Fanaroff}) constitute the parent population of FSRQs, while low-power radio galaxies of the Fanaroff-Riley morphological type I (FR\\,Is) are considered to be ``misaligned'' BL Lacs (\\citealt{1989ApJ...336..606Barthel}; \\citealt{1995PASP..107..803Urry}; \\citealt{2009ApJ...694L.107Xu}).  High-dynamic range radio imaging of BL Lac objects at GHz frequencies often reveals the presence of a diffuse extended radio emission, with the integrated luminosity exceeding, in several cases, the FR\\,I/FR\\,II divide. We note that this borderline scales with the optical luminosity of the host as $\\propto L_{\\rm opt}^{1.8}$, with the normalization $L_{\\rm 1.4\\,GHz} = 10^{24}$\\,W\\,Hz$^{-1}$ for the absolute isophotal magnitude of the host galaxy $-21$ measured to 24.5 magnitudes per arcsec$^2$ in the rest-frame of a source (\\citealt{1996AJ....112....9Ledlow}). This may indicate that the parent population of BL Lac objects is a mixture of FR\\,I and FR\\,II type galaxies (see, e.g., \\citealt{1999A&AS..139..601Cassaro}; \\citealt{2001AJ....122..565Rector}; \\citealt{2008MNRAS.391..967Landt}; \\citealt{2010ApJ...710..764Kharb}). Therefore, a careful investigation of the extended diffuse radio emission in BL Lacs is of importance for the models relating the AGN accretion power and jet luminosity in general, and the blazar population studies in particular (e.g., \\citealt{1982MNRAS.199..883Blandford}; \\citealt{1991Natur.349..138Rawlings}; \\citealt{2011MNRAS.414.2674Ghisellini}).  The high frequency-peaked BL Lacs (HBLs), i.e. blazars with synchrotron peak frequencies $> 10^{15}$\\,Hz (see the discussion in \\citealt{2010ApJ...716...30Abdo} and references therein), dominate the population of extragalactic TeV emitters\\footnote{\\texttt{http://tevcat.uchicago.edu/}}. These sources, unlike the low frequency-peaked BL Lacs (LBLs; synchrotron peak frequencies $<10^{15}$\\,Hz), seem to be associated strictly with FR\\,I-type large-scale radio structures (e.g., \\citealt{2010ApJ...710..764Kharb}), and to form a relatively distinct sub-class of blazars with respect to the jet physical properties (e.g., \\citealt{2010MNRAS.401.1570Tavecchio}). Also, no superluminal velocities have been detected in the TeV-emitting HBLs on milli-arcsec scales till now, despite several dedicated observational programs (e.g., \\citealt{2010ApJ...723.1150Piner} and references therein). Such apparent superluminal velocities are typical for FSRQs and not rare for LBLs.  SHBL\\,J001355.9$-$185406 (J2000.0 R.A.\\,$\\rm=00^{h}13^{m}56\\fs054$, Dec.\\,$\\rm=-18\\degr54\\arcmin06\\farcs48$; redshift $z = 0.0948$ according to \\citealt{2009MNRAS.399..683Jones}) is a peculiar example of an HBL. This source was listed in a catalog of extremely high X-ray--to--radio flux ratio targets in the multi-frequency `Sedentary Survey' of HBLs performed by \\citet{2005A&A...434..385Giommi}. Its $0.1-2.4$\\,keV X-ray flux is $1.26\\times10^{-11}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$, and the corresponding $1400$\\,MHz radio flux spectral density is $29.2\\pm1.0$\\,mJy (see, \\citealt{1998AJ....115.1693Condon}). Very recently, this source has been detected at the level of $\\sim 1\\%$ of the flux of the Crab nebula above 300\\,GeV by the High Energy Stereoscopic System (H.E.S.S.; \\citealt{2013A&A...554A..72Hess}). The analysis of the {\\it Fermi} Large Area Telescope ({\\it Fermi}-LAT) data revealed a faint and flat-spectrum counterpart of SHBL\\,J001355.9$-$185406 at the GeV photon energies, with the integrated photon flux above $100$\\,MeV of $(0.9\\pm0.7)\\times10^{9}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ and the photon index of $1.5\\pm0.2$ (\\citealt{2010ATel.3014....1Sanchez}). At lower frequencies, the discussed object was also reported as a bright and variable NIR/optical emitter. In particular, the NIR flux of SHBL\\,J001355.9$-$185406 determined in 2010 was about 0.7\\,mag brighter than that previously published in the {\\emph {2MASS}} survey (\\citealt{2010ATel.3023....1Cassarco}). The corresponding $70\\%$ flux increase imply the jet origin of the NIR continuum.  \\begin{figure}[t] \\includegraphics[height=8.0cm,width=8.0cm]{n.jpg} \\caption{ Contour plot of the 1400\\,MHz {\\emph {NVSS}} radio continuum emission. The plus sign marks the position of blazar core and the bottom left inset shows the size of restoring beam. } \\label{fig.nvss} \\end{figure}  The peculiarity of the SHBL\\,J001355.9$-$185406 is related however to its large-scale radio structure which have been found when examining the archival radio National Radio Astronomy Observatory ({\\emph {NRAO}}) Very Large Array ({\\emph {VLA}}) Sky Survey ({\\emph {NVSS}}; \\citealt{1998AJ....115.1693Condon}) maps of the TeV-detected BL Lacs. This structure, shown in Figure\\,\\ref{fig.nvss}, consists of a compact core component, as expected for a blazar, but also of a one-sided and elongated lobe/jet-like feature extending to the south-west from the nucleus. At the redshift of the source (luminosity distance of $d_{L} = 430$\\,Mpc for the flat cosmology with $H_{0} = 71$ km sec$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} =0.27$ and $\\Omega_{\\Lambda} =0.73$), the 3\\farcm3 extension of the lobe corresponds to the physical (projected) size of $\\sim 344$\\,kpc and the $1400$\\,MHz flux $15.8$\\,mJy corresponds to the lobe's monochromatic luminosity of $\\sim 5 \\times 10^{39}$\\,erg\\,s$^{-1}$. The derived extended luminosity is therefore comfortably low as for the low-power blazar of the HBL type, yet the morphology of the giant lobe is striking for two reasons.  First, the lobe is clearly edge-brightened, and hence the whole large-scale radio structure of SHBL\\,J001355.9$-$185406 resembles more an FR\\,II-type object. Second, the extended structure in the discussed source is clearly one-sided. This is \\emph{not} expected for the large-scale structures in radio-loud AGN in general, regardless on their Fanaroff-Riley type, since the lobes' expansion velocities on scales of tens and hundreds of kpc are expected to be sub-relativistic only. In other words, the Doppler (and time-travel) effects cannot be easily blamed for the apparent lack of the counter-lobe in SHBL\\,J001355.9$-$185406 (but see in this context \\citealt{2004ApJ...613..119Stawarz} and \\citealt{2007ApJ...670L..85Bagchi}). In order to investigate this issue in more detail, we performed deep and high-resolution radio observations of SHBL\\,J001355.9$-$185406 at low frequencies. The results of the analysis are presented below.  \\begin{figure}[t] \\includegraphics[height=6.5cm,width=6.5cm]{spec_peak.jpg} \\caption{ Integrated radio spectrum of the unresolved core in SHBL\\,J001355.9$-$185406. The dashed line shows the result of weighted linear regression analysis. } \\label{spec_radio} \\end{figure} ",
        "conclusions": "\\label{conclutions}  In our study presented here, we investigated in detail large-scale radio morphology of the TeV-detected blazar SHBL\\,J001356$-$185406 which, on the low-resolution {\\emph {NVSS}} map, appears to be of the FR\\,II type. High-sensitivity and high-dynamic range (Table\\,\\ref{obs}) observations (which do not suffer from a missing flux; see Section\\,\\ref{result}) were made using the {\\emph {GMRT}} at low frequencies of 156, 259 and 629\\,MHz. Our analysis indicated that: (1) the integrated radio spectral index is flat down to 156\\,MHz for the blazar core (Figure\\,\\ref{spec_radio}); (2) no diffuse arcmin-scale emission is present around the unresolved blazar core (see Figures\\,\\ref{map.150}--\\ref{map.610}); (3) the 3\\farcm3-long, lobe/jet-like structure extending to the south-west from the blazar core on the {\\emph {NVSS}} map breaks into three distinct, steep-spectrum radio features, unrelated to the blazar but instead associated most likely with background AGN.   As a corollary to this study we point out that the observations presented here are rather unique, since very low-frequency, high-resolution and high-dynamic range studies of TeV detected HBLs are particularly sparse. The new data allow us to constrain the upper limits for the extended radio halo around the blazar as $< 12\\%$ at 156\\,MHz, $< 8\\%$ at 259\\,MHz, and $< 0.9\\%$ at 629\\,MHz of the core emission. This is agreement with the early {\\emph {VLA}} studies of large-scale structures of BL Lacs at GHz frequencies by \\citet{1984AJ.....89..189Ulvestad} and \\citet{1986AJ.....92....6Ulvestad}, who found that only a few of the objects from the analyzed samples possess any arcminute-scale structures with total fluxes exceeding $\\sim 0.1\\%$ of the core emission.  The another interesting result of our study is that the radio spectral index of the SHBL\\,J001356$-$185406 core remains flat down to 156\\,MHz, with no dramatic curvature. Flat spectra (indices $\\alpha \\sim 0$) of relativistic jets observed at small viewing angles are widely believed to arise due to a superposition of self-absorbed spectra of different parts of an outflow with conserved magnetic flux and maintained power-law spectrum of the emitting electrons \\citep{1979ApJ...232...34Blandford}. The fact that in the case of the SHBL\\,J001356$-$185406 core we do not see any spectral turnover down to 156\\,MHz indicates therefore that a self-similar character of the jet in the source holds up to relatively large distances from the jet base."
    },
    "1402/1402.3325_arXiv.txt": {
        "abstract": "We developed a Keplerian-based Hamiltonian splitting for solving the gravitational $N$-body problem. This splitting allows us to approximate the solution of a general $N$-body problem by a composition of multiple, independently evolved $2$-body problems. While the Hamiltonian splitting is exact, we show that the composition of independent $2$-body problems results in a non-symplectic non-time-symmetric first-order map. A time-symmetric second-order map is then constructed by composing this basic first-order map with its self-adjoint. The resulting method is precise for each individual $2$-body solution and produces quick and accurate results for near-Keplerian $N$-body systems, like planetary systems or a cluster of stars that orbit a supermassive black hole. The method is also suitable for integration of $N$-body systems with intrinsic hierarchies, like a star cluster with primordial binaries. The superposition of Kepler solutions for each pair of particles makes the method excellently suited for parallel computing; we achieve $\\gtrsim 64\\%$ efficiency for only $8$ particles per core, but close to perfect scaling for $16384$ particles on a $128$ core distributed-memory computer. We present several implementations in \\texttt{Sakura}, one of which is publicly available via the AMUSE framework. ",
        "introduction": "\\label{sec1:introduction}  Since the pioneering work of \\cite{vonHoerner60}, \\cite{Aarseth63} and \\cite{vanAlbada68} $N$-body simulations have been an essential tool for the theoretical understanding of self-gravitating astrophysical systems. Such systems often show a large dynamic range of time-scales. Thus, instead of a fixed or adaptive global time-step, most of the $N$-body codes adopt individual or block time-step algorithms in order to advance the particles in time \\citep{McMillan86, MakinoAarseth92, Aarseth03}. In addition, different approaches to calculate the acceleration of each particle, such as using grids \\citep{HockneyEastwood88} or a hierarchical tree data structure \\citep{BarnesHut86}, are commonly employed to decrease the computational cost of the simulations. These approaches allow the use of a larger number of particles, despite only giving an approximation to the true acceleration of each particle. Therefore, these codes should not inadvertently be used in simulations of collisional systems such as planetary systems, dense star clusters or the inner parts of galactic nuclei.  In collisional systems the individual interactions between particles play an important role in the dynamical evolution of the system as a whole. For example, the formation of hard binaries in star cluster core collapse \\citep{Tanikawa12} requires very precise integration methods to correctly evolve close encounters between particles. This precision is only possible if we use more accurate, direct brute-force methods, to calculate the accelerations due to each pair of particles in the system. The main difficulty here is that with the formation of the first hard binary in the system, the simulation as a whole experiences a slow-down in performance due to the necessity to decrease the time-step size in order to accurately integrate such compact sub-systems.  Currently, the  most effective and common approach to overcome such obstacles seems to be a combination of the block time-step algorithm \\citep{McMillan86}, Ahmad-Cohen neighbour scheme \\citep{AhmadCohen73} and some sort of $2$-body regularization \\citep{PretoTremaine1999, MikkolaTanikawa1999a, MikkolaTanikawa1999b, MikkolaAarseth2002} in order to handle very compact sub-systems efficiently. This is the approach used in modern Hermite integrators for collisional stellar systems \\citep{MakinoAarseth92, Aarseth03}.  In this paper, we develop a new Keplerian-based Hamiltonian splitting for the gravitational $N$-body problem. This splitting allows us to approximate the solution of a general $N$-body problem by a composition of independently evolved $2$-body problems. While the Hamiltonian splitting is exact, we show in section~\\ref{sec2:method} that the composition of independent $2$-body problems results in a non-symplectic non-time-symmetric first-order map. A time-symmetric second-order map is then constructed by composing this basic first-order map with its self-adjoint. The advantages of this Keplerian-based integrator are: i) a guarantee that every pair of particles is always integrated precisely; ii) the method does not suffer from slow-down in performance when tight binaries are present in the simulation, and iii) the method allows for good parallel efficiency \\citep{PortegiesZwart2001}. ",
        "conclusions": "\\label{sec5:discussion}  We have described a Keplerian-based Hamiltonian splitting for gravitational $N$-body simulations and its implementation in a new code called \\texttt{Sakura}. In this method a general $N$-body problem can be solved as a composition of multiple, independent, $2$-body problems. The integration scheme is constructed on the assumption that, during a small time interval $\\tau$, each pair of particles in the system can be treated as an independent $2$-body problem. With this splitting an analytical \\texttt{kepler\\_solver} can be used to accurately, and independently, evolve each $2$-body interaction in the system, thus making the code especially suitable for simulations in which compact primordial binaries or close dynamically formed binaries are present. Hierarchies in which one of the components is a compact binary and systems with a central dominant mass are also examples of physical systems in which \\texttt{Sakura} performs well when compared to traditional codes.  Because \\texttt{Sakura} can easily handle arbitrarily compact binaries in an $N$-body simulation, the code is able to evolve a star-cluster through core-collapse without much difficulty. In particular, since \\texttt{Sakura} can do this even with the use of constant time-steps, the simulation does not suffer from any slow down in performance as is the case in other non-regularized $N$-body codes. As an example, in the $1024$-bodies core-collapse simulation presented in section~\\ref{sec3.2:largen}, \\texttt{Sakura} was able to complete the run in about $3$ days of CPU time on a $4$-core machine. The same system being integrated with a $4$-th order Hermite integrator took more than one month of CPU time on the same machine, due to a severe slow down in performance after the formation of the first hard binary in the system.  There are, however, some circumstances in which \\texttt{Sakura} may not be the most suitable code to perform an $N$-body simulation. For example, for systems in which multiple bodies democratically interact among themselves, \\texttt{Sakura} may perform almost as badly as a simple Leapfrog integrator, as demonstrated in the integrations of a figure-eight system in section~\\ref{sec3.1:smalln}. This happens because of our underlying assumption that the $N$-body problem can be decomposed in multiple, independent, $2$-body problems. Such decomposition in fact constitutes the main source of error when a given $i-j$ pair is being integrated with a time-step $\\tau$ which is larger than the time-scale of the perturbation due to a close neighbour. In many cases this issue may be surpassed by decreasing the constant time-step size used in the simulation. However, the cause of the problem lies on the non-commutativity of $2$-body interactions when multiple bodies are involved in a democratic close encounter. While it is not easy to solve this issue without breaking our Keplerian splitting approach, the introduction of an adaptive time-step scheme in \\texttt{Sakura} might alleviate these numerical difficulties and is currently under investigation.  According to some of our tests (not reported in the present paper), a time-step criterion based on the strength of the perturbation on a given $i-j$ pair seems to work relatively well compared to a constant $\\tau$. However, this improvement is only significant when close multiple-body encounters take place. On the other hand, one could in principle choose $\\tau \\sim \\min(r_{ij}/v_{ij})$, $\\tau \\sim \\min((r_{ij}/a_{ij})^{1/2})$ or use a traditional Aarseth-like time-step criterion, but we advocate that this may not be the optimal choice because these criteria also include the contribution of the $i-j$ pair itself, which in principle contributes to a severe decrease in time-steps if a close binary is present in the system. In \\texttt{Sakura}, these severely short time-steps are not necessary, because the use of a Keplerinan treatment for each pair of particles automatically regularizes every $2$-body interaction in the system. It is only when multiple-body encounters happens that the time-step should adapt itself to properly resolve the approximation of a perturber. Therefore, we stress here our preference for a perturbation-based time-step criterion rather than an Aarseth-like criterion for use in \\texttt{Sakura}. Whether or not such perturbation-based criterion is the best choice for \\texttt{Sakura} is a matter that will be addressed elsewhere.  Another point we want to emphasize here is the behaviour of \\texttt{Sakura} when integrating a system with a central massive black-hole. As shown in Fig.~\\ref{fig:03}, the level of energy conservation in Leapfrog and $4$-th order Hermite integrations remains approximately constant with the increase of the black-hole to stellar mass ratio. For \\texttt{Sakura}, we found that it performs much better than previous approaches, becoming gradually more precise with the increase of the mass ratio. In particular, for the case of a mass ratio $q = 10^{6}$ \\texttt{Sakura} can give $\\gtrsim 5$ orders of magnitude better energy conservation than Hermite integrator, being at the same time up to $4$ orders of magnitude faster when the mass ratio increases to $q\\gtrsim 10^{9}$. The fact that \\texttt{Sakura} can be, at the same time, fast and accurate in this regime, makes this code highly suitable for nearly Keplerian systems where a massive particle dominates the evolution of surrounding particles, such as in planetary systems and galactic nuclei with super-massive black-holes.  Lastly, \\texttt{Sakura} has proven to be quite easy to parallelize for distributed memory systems using \\texttt{MPI}. The GPU implementation, even though theoretically easy, is still not totally efficient due to the presence of many branching conditions in the Kepler-solver. In algorithmic terms, the bulk of computation in \\texttt{Sakura} occurs inside a double loop, similar to the one used to calculate the acceleration of particles in conventional $N$-body codes. Therefore, we were able to immediately employ existent parallelization schemes in \\texttt{Sakura} without much effort. We argue that the fact that our GPU implementation is not yet very efficient is not a problem due to the paralelization scheme itself, but rather due to the poor/inefficient support for branch conditions in current GPUs. A restructure in our Kepler-solver in order to eliminate (or minimize) these branch conditions may address this issue, and possibly speed up even more the \\texttt{MPI} version on CPUs, which has already shown a remarkable parallel efficiency, with close to $100\\%$ efficiency for $16\\tilde{k}$ particles on $128$ cores, and $64\\%$ efficiency when using only $8$ particles per core."
    },
    "1402/1402.0326_arXiv.txt": {
        "abstract": "{The \\textit{Herschel} Reference Survey is a complete volume-limited, K-band-selected sample of nearby objects including Virgo cluster and isolated objects. Using a recent compilation of H{\\sc i} and CO data for this sample we study the effects of the cluster environment on the molecular gas content of spiral galaxies. With the subsample of unperturbed field galaxies, we first identify the stellar mass as the scaling variable that traces the total molecular gas mass of galaxies better. We show that, on average, H{\\sc i}-deficient galaxies are significantly offset (4 $\\sigma$) from the $M(H_2)$ vs. $M_{star}$ relation for H{\\sc i}-normal galaxies. We use the $M(H_2)$ vs. $M_{star}$ scaling relation to define the H$_2$-deficiency parameter as the difference, on logarithmic scale, between the expected and observed molecular gas mass for a galaxy of given stellar mass. The H$_2$-deficiency parameter shows a weak and scattered relation with the H{\\sc i}-deficiency parameter, here taken as a proxy for galaxy interactions with the surrounding cluster environment. We also show that, as for the atomic gas, the extent of the molecular disc decreases with increasing H{\\sc i}-deficiency. All together, these results show that cluster galaxies have, on average, a lower molecular gas content than similar objects in the field. Our analysis indicates that ram pressure stripping is the physical process responsible for this molecular gas deficiency. The slope of the $H_2-def$ vs. $\\mathrm{\\hi}-def$ relation is less than unity, while the $D(\\mathrm{\\hi})/D(i)$ vs. $\\mathrm{\\hi}-def$ relation is steeper than the $D(CO)/D(i)$ vs. $\\mathrm{\\hi}-def$ relation, thereby indicating that the molecular gas is removed less efficiently than the atomic gas. This result can be understood if the atomic gas is distributed on a relatively flat disc that is more extended than the stellar disc. It is thus less anchored to the gravitational potential well of the galaxy than the molecular gas phase, which is distributed on an exponential disc with a scalelength $r_{CO}$ $\\simeq$ $0.2 r_{24.5}(g)$. There is a clear trend between the $NUV-i$ colour index, which is a proxy for the specific star formation activity, and the H$_2$-deficiency parameter, which suggests that molecular gas removal quenches the activity of star formation. This causes galaxies migrate from the blue cloud to the green valley and, eventually, to the red sequence. The total gas-consumption timescale of gas deficient cluster galaxies is comparable to that of isolated, unperturbed systems. The total gas depletion timescale determined by considering the recycled fraction is $\\tau_{gas,R}$ $\\simeq$ 3.0-3.3 Gyr, which is significantly larger than the typical timescale for total gas removal in a ram pressure stripping process, indicated by recent hydrodynamical simulations to be $\\tau_{RP}$ $\\simeq$ 1.5 Gyr. The comparison of these timescales suggests that ram pressure, rather than a simple stop of the infall of pristine gas from the halo, will be the dominant process driving the future evolution of these cluster galaxies.   } {} {} {} {} {} ",
        "introduction": "The molecular gas phase plays a fundamental role in the process of star formation in galaxies. The atomic gas, which is dominant in normal, field late-type galaxies, collapses within molecular clouds to form new stars. This simple matter cycle can be significantly perturbed in dense environments (e.g. Boselli 2011). The atomic gas phase, which in normal, unperturbed galaxies extends up to $\\sim$ 1.8 times the stellar disc (Cayatte et al. 1994; Bigiel \\& Blitz 2012), can be easily removed during the interaction of galaxies with the surrounding environment (Boselli \\& Gavazzi 2006). Since the seminal work of Haynes \\& Giovanelli (1984), there has indeed been strong observational evidence that late-type cluster galaxies have less atomic gas than similar objects in the field (e.g. Cayatte et al. 1990; Solanes et al. 2001; Gavazzi et al. 2005). Tidal interactions with nearby companions or with the cluster gravitational potential well (galaxy harassment, Moore et al. 1998), the dynamical interaction between the hot intracluster medium (ICM) with the cold interstellar medium (ISM) of galaxies moving at high velocity ($\\sim$ 1000 km s$^{-1}$) within clusters (ram pressure stripping, Gunn \\& Gott 1972; thermal evaporation, Cowie \\& Songaila 1977) can remove the gaseous component.  The lack of gas naturally quenches star formation, thereby transforming active systems in quiescent objects (e.g. Boselli \\& Gavazzi 2006). The timescales relative to this transforming process, and the structural, physical and kinematical properties of the perturbed galaxies drastically change according to the perturbing process. Gravitational interactions indistinctly perturb all galaxy components (dark matter, stars, gas, dust, etc.) thereby producing strong signatures in the morphological and kinematical structure of the perturbed objects. To be efficient in fully transforming the whole late-type galaxy population falling into rich clusters, however, they require very long timescales since multiple encounters are necessary. Indeed, given the high-velocity dispersion within clusters ($\\sim$ 1000 km s$^{-1}$), the perturbative phase during a fly-by encounter is too short to induce strong effects (Boselli \\& Gavazzi 2006). The dynamical interactions with the hot, dense ICM and, in particular, the ram pressure stripping phenomenon, has been identified as the dominant process perturbing late-type galaxies in nearby rich clusters (e.g. Vollmer et al. 2001; Boselli \\& Gavazzi 2006; Tonnesen et al. 2007). Observations, models and simulations all indicate that the atomic gas phase can be easily perturbed on relatively short timescales ($\\sim$ 100-200 Myr; Vollmer et al. 2004; Roediger et al. 2005; Boselli et al. 2006; 2008a, 2008b; Crowl \\& Kenney 2008). The effects on the molecular gas phase are still under study. Modelling the effects of ram pressure on a multiphase gas disc is still challenging since it requires taking several physical effects into account, such as heating and cooling of the ISM, self-gravity, star formation, stellar feedback, and magnetic field. Given the fractal distribution of the gaseous component within the ISM, models must also simultaneously predict the gas distribution on different scales, from giant molecular clouds to the tails whose size can exceed the size of galaxies (Roediger 2009). Recent attempts to simulate the effects of ram pressure on a multiphase ISM on different scales have been carried out by Tonnesen \\& Bryan (2009). These authors have shown that the molecular gas phase, even in high-density regions, can be removed during the interaction. Whereas detailed multifrequency observations of representative objects in nearby clusters begin to show cases of ongoing molecular gas ram pressure stripping events (Vollmer et al. 2008a, 2009, 2012a; Sivanandam et al. 2010), strong and convincing statistical evidence is still lacking. Since star formation occurs mainly inside molecular clouds, the perturbation of the molecular gas phase in any kind of interaction can have a strong impact on the evolution of the stellar populations of galaxies. It is thus paramount to understand whether the molecular hydrogen is perturbed or not and removed in high-density environments.\\\\   The first studies of the molecular gas content of cluster galaxies and the search for any possible evidence of molecular gas stripping come from Kenney \\& Young (1989) for the Virgo cluster and from Casoli et al. (1991) for the Coma cluster. Kenney \\& Young (1988; 1989) discovered low-luminosity Virgo cluster spirals characterised by weak CO emission lines, which they interpreted as the first proof of a low molecular hydrogen content in this cluster. Casoli et al. (1991), on the other hand, observed normal molecular gas content in a few spiral galaxies in the Coma cluster and deduced that the cluster environment is not able to remove the molecular phase, which is strongly anchored to the deep gravitational potential well of galaxies. These results were later questioned by Boselli et al. (1997) as due to strong biases. Boselli et al. (1997) have shown how the low CO emission of the low-luminosity Virgo cluster galaxies discovered by Kenney \\& Young (1988) was not evidence of the lack of H$_2$, but rather an effect related to the decrease in the conversion factor $X_{CO}$ in low-luminosity, metal-poor objects with respect to massive, metal-rich galaxies. Indeed, in these low-metallicity systems the interstellar radiation field is able to dissociate the molecular gas in the diffuse medium, where $X_{CO}$ is low, thus increasing the mean $X_{CO}$ factor of these galaxies with respect to massive objects (Boselli et al. 1995, 1997, 2002). The result of Casoli et al. (1991) has also been criticised because the observed Coma cluster galaxies were selected in the far-infrared. Given the tight correlation between the far-infrared and the CO emission, this selection criterion might have excluded molecular gas poor objects from the observed sample (Boselli et al. 1997, 2002). Other CO surveys of nearby clusters have been published in Lavezzi \\& Dickey (1998) and in Scott et al. (2013). Using a collection of CO data taken from the literature for a large sample of nearby galaxies, Jablonka et al. (2013) have recently shown that, on average, cluster galaxies have a lower CO luminosity per unit far-infrared luminosity than field objects. They interpreted this result as an evidence of molecular gas stripping in high-density environments. This result, however, is based on a heterogeneous sample of isolated and cluster galaxies often selected according to different criteria and thus might be biased by subtile and complex selection effects (e.g. Boselli et al. 1997; 2013c). The lack of cluster and field samples selected according to similar criteria prevented previous works from deriving any firm conclusion on this topic. In recent years, several efforts have been undertaken to construct a sample of isolated, unperturbed galaxies with molecular gas data that are ideally defined to be used as references for any statistical study. Sauty et al. (2003) made systematic CO observations of galaxies in the catalogue of isolated objects constructed by Karachentseva (1973). Observations of these isolated objects were later extended by Lisenfeld et al. (2011). However, these new datasets have not yet been used to quantify the molecular gas content of nearby cluster galaxies.  In the preparation of the \\textit{Herschel} mission (Pilbratt et al. 2010) we began a \\textit{SPIRE} guaranteed time key project aimed at studying the properties of the ISM of a K-band-selected, volume-limited sample of 322 nearby galaxies. This project, named the \\textit{Herschel} Reference Survey (HRS), has been extensively described in Boselli et al. (2010). To provide the community with the largest possible set of multifrequency data necessary for any kind of statistical analysis, we have collected data at all wavelengths, from UV (Boselli et al. 2011; Cortese et al. 2012a) to radio centimetric, including H$\\alpha$ narrow-band imaging (Boselli et al. in prep.), integrated optical spectroscopy (Boselli et al. 2013a), PACS (Cortese et al. 2014) and SPIRE photometry (Ciesla et al. 2012). For such purpose, we have recently obtained new CO data for a large fraction of the late-type galaxies of the HRS. We combined these data with those available in the literature and produced a homogenised catalogue of $^{12}$CO(1-0) data for 225 galaxies of the HRS sample. We also collected from the literature and homogenised H{\\sc i} data for 315 out of the 322 HRS galaxies (256/260 for spirals; Boselli et al. 2013b, Paper I). The quality of the data at hand, the accurate estimate of the total CO flux emission and of their uncertainties, and the presence of a well defined, complete sample including cluster and isolated galaxies selected according to similar criteria, all allow us to extend and revisit the critical topic of the effects of the environment on the molecular gas properties of cluster galaxies on the best statistically significant sample with CO data available in the literature. The study of the effects of the environment on the atomic gas content of cluster galaxies has been presented in Cortese et al. (2011), the one on the dust mass in Cortese et al. (2010, 2012). In an accompanying paper (Boselli et al. 2013c; paper II of this series), we study the molecular and total gas scaling relations of the HRS. ",
        "conclusions": "We use the recent compilation of homogeneous atomic and molecular gas data for the \\textit{Herschel} Reference Survey, a K-band-selected, volume-limited, complete sample of galaxies in the nearby universe, to study the effects of the surrounding environment on the molecular gas content of cluster galaxies. Using the subsample of isolated, unperturbed objects, we first identify the $M(H_2)$ vs. $M_{star}$ relation as the less dispersed scaling relation necessary to estimate the expected molecular gas content of perturbed galaxies of different size and luminosity. By studying the distribution of the distance to this relation we show that H{\\sc i}-deficient, Virgo cluster galaxies have, on average, less molecular gas than H{\\sc i}-normal field objects. We then calibrate the H$_2$-deficiency parameter, defined as the difference between the expected and the observed molecular gas mass of an unperturbed galaxy of similar stellar mass. The H$_2$-deficiency parameter increases with H{\\sc i}-deficiency parameter, here taken as tracer of an undergoing perturbation. The relation is weak and scattered, but statistically significant (the probability that the two variables are correlated is $P$ $>$ 99 \\%). This relation is barely detectable in the whole sample because of the large uncertainty on the total CO fluxes due to aperture corrections. It is, however, confirmed on a subsample of galaxies with high-quality CO data from mapping. Using a subsample of seventeen CO mapped galaxies, we show that, as for the atomic gas, the extension of the molecular disc decreases with increasing H{\\sc i}-deficiency. This result suggests that gas removal is an outside-in process able to truncate the gaseous disc of galaxies in high-density regions.  The analysis presented in this work shows the presence of molecular gas deficient galaxies in the core of the Virgo cluster. It is consistent with previous results based on the study of the radial distribution of the molecular gas of a few, resolved galaxies (Boselli et al. 2006; Fumagalli et al. 2009) and on the statistical analysis of a heterogeneous sample of cluster and isolated objects (Jablonka et al. 2013). This result is statistically significant, based on a homogeneous sample of cluster and field galaxies selected according to similar criteria to minimise selection biases, and robust vs. the adoption of a constant or a luminosity-dependent conversion factor. Combined with the recent evidence of dust stripping obtained from \\textit{Herschel} observations (Cortese et al. 2010, 2012b), the deficiency of molecular gas in cluster galaxies can be interpreted as an effect of the ram pressure stripping exerted by the hot intergalactic medium on galaxies moving at high velocity within the cluster rather than due to a low formation rate of molecular gas whenever the interstellar medium of galaxies drops below a critical density, as firstly proposed by Fumagalli et al. (2009). Given the steep gradient in the molecular gas distribution over the galactic disc, molecular gas stripping is less efficient than atomic gas stripping, as indicated by the slope $<$ 1 in the $H_2-def$ vs. $\\mathrm{\\hi}-def$ relation and by the steeper slope in the $D(\\mathrm{\\hi})/D(i)$ vs. $\\mathrm{\\hi}-def$ (slope = -0.50) than in the $D(CO)/D(i)$ vs. $\\mathrm{\\hi}-def$ (slope = -0.36) relation. The most H{\\sc i}-deficient, Virgo cluster galaxies ($\\mathrm{\\hi}-def$ $\\gtrsim$ 1) have, on average, a factor of $\\sim$ 2 less molecular gas than similar objects in the filed.  The molecular and total gas consumption timescales of these H$_2$-deficient cluster galaxies are $\\tau_{H_2}$ $\\simeq$ 1.2 Gyr, $\\tau_{gas}$ $\\simeq$ 2.6 Gyr, and $\\tau_{gas,R}$ $\\simeq$ 3.0-3.3 Gyr when the recycled gas is considered. These timescales are significantly longer than the timescale necessary for ram pressure to produce a typical H{\\sc i}-deficient cluster galaxy ($\\sim$ 0.1 Gyr). On these timescales a galaxy can cross the core of the cluster and thus reach the peak of its gas stripping $\\simeq$ three times. The total gas depletion timescale determined considering the contribution of the recycled gas produced by stars after the first crossing of the cluster is also a factor of $\\gtrsim$ 2 longer than the typical timescale for total gas ablation via ram pressure determined by recent hydrodynamical simulations to be of the order of $\\lesssim$ 1.5 Gyr. This indicates that ram pressure stripping, rather than gas consumption via star formation, is the main process driving the future evolution of these cluster gas-deficient spiral galaxies. The removal of the cold gaseous phase quenches the activity of star formation, inducing the migration of these originally star forming systems from the blue sequence into the green valley and, eventually into the red sequence."
    },
    "1402/1402.0783_arXiv.txt": {
        "abstract": "This paper aims to investigate/map the effects that perturbations applied to an accretion disk might produce on the registered Light Curves (LC). The case of accretion disks around supermassive active black holes (AGNs) is studied with the goal to explain some of the statistical properties of the observed IntraDay Variability (IDV). The region producing optical IDV is perturbed by allowing it to develop a mass density of a fractional Brownian Motion-like type. The light curves and spectral slopes are calculated and compared to observational data for different Hurst parameters. The spectral slopes of the simulated light curves vary in the range $(0.4,2.5)$. The agreement with observational data shows that a magnetized disk subjected to stochastic perturbations can produce some of the features observed in the light curves.\\\\ \\textbf{Keywords}: accretion, accretion discs; magnetohydrodynamics (MHD); fractional Brownian Motion ",
        "introduction": "\\label{sect:Perturbation}  The equilibrium state is perturbed by a prescribed temporal variation in the density. When the perturbation is applied, the new physical parameters that satisfy the MHD equations consist of a sum of two parts: the equilibrium part and the perturbed part.  The perturbations are isothermal. This is not because we believe this is necessarily the case in an astrophysical disk. We make this assumption because in this toy model we want to study fluctuations in energy occurring at the same temperature, and not the distribution of energy in a large interval of temperatures.  Accretion is possible due to the outward angular momentum transport by a sum of Reynolds and Maxwell stresses~\\cite{art:21b,art:11b,art:18b} and we assume that the perturbations in this mechanism are responsible for the variability. The component of the stress tensor responsible for the luminosity is defined as  \\begin{equation} m_{r\\phi} (r,t) = \\mu _0 \\rho _0 (r) v_{r} (r,t) v_{\\phi} (r,t) - B_{r} (r,t) B_{\\phi} (r,t), \\end{equation} where $v_r, v_\\phi, b_r, b_\\phi$ are the perturbed components of the velocity and magnetic field.  The luminosity emitted from a patch of the disk due to the perturbation, between radii $r_1$ and $r_2$ is~\\cite{art:21b}  \\begin{equation} L_{12} (t) = \\int _{r_1} ^{r_2} m_{r\\phi} (r,t) dr \\end{equation} and this equation is the starting point of producing the light curves shown in the Results section of the paper.  \\subsection{Stochastic density perturbation\\label{sect:densFBM}}  There is a general consensus that the underlying mechanism producing IDV is stochastic~\\cite{art:29l,art:16a,art:50m}.  The time dependency of the density perturbation is taken to be the realization of a stochastic process such that $\\frac{\\partial \\rho _1}{\\partial t} = A_\\rho B_H (t) = \\xi$, where $B_H (t)$ is a normalized fractional Brownian Motion (fBM) and $A_\\rho$ is a constant amplitude, which will be prescribed by us. It is expected that locally fluctuating component of magnetic field create numerous current sheets inside the disk~\\cite{art:17m}, which can be used as random reconnection sites. The accretion disk can work as a dynamo amplifying the stochastic magnetic field~\\cite{book:22b}. This happens because even in a perfectly conducting plasma, turbulence develops through the MagnetoRotational Instability (MRI)~\\cite{art:42s}. The already mentioned peculiar shape of the observational PSD prompts us to consider that the perturbation is a signal with non-Markovian properties. A sample light curve is shown in Fig.~\\ref{fig:idv}.  There is no constraint imposed to the value of $H$ in our framework mainly because there is a lack of observational data to correlate IDV characteristics to characteristics of the central engine. For example, \\cite{art:7n} report, regarding the optical IDV  of S5 0716+714, that there is no correlation between the source magnitude and the amplitude of the IDV and that there is no clear correlation between the source magnitude and the rate of magnitude variation; \\cite{art:5w} report that time scales of optical variability in BL Lacs do not correlate with luminosity; \\cite{art:50m} conclude that it cannot be assumed that quasars with similar mass will necessarily have the same variability properties; \\cite{art:7h} discuss optical data of a sample of Bl Lac and state that there are no dependencies of the time scales on intrinsic properties, redshift or absolute magnitude.  \\begin{figure} \\begin{center} \\includegraphics*[width=6cm,angle=0]{B1LC.eps} \\end{center} \\caption{B (blue) band time series for the object Bl Lacertae S5 0716+714; JD stands for Julian Day.\\label{fig:idv}} \\end{figure}  The \\emph{definition} of fBM is~\\cite{art:52m}  \\begin{eqnarray} \\lefteqn{B_H(t) = \\frac{1}{\\Gamma \\left ( H + \\frac{1}{2}\\right)}\\left( \\int _{-\\infty}^0 \\left [ (t-s)^{H-1/2}-(-s)^{H-1/2}\\right ] dB(s) \\right)}\\nonumber\\\\ && +\\frac{1}{\\Gamma \\left ( H + \\frac{1}{2}\\right)}\\left(\\int _0^t (t-s)^{H-1/2}dB(s)\\right), \\end{eqnarray} where $\\Gamma$ represents the special function  \\begin{equation} \\Gamma(z) = \\int _{0}^{\\infty} x^{z -1}e^{-x}dx. \\\\ \\end{equation}  The Power Spectral Distribution of such a signal depends on the Hurst parameter $H$ as~\\cite{art:52m}  \\begin{equation} P(f)\\sim f^{1-2H}.\\label{eq:psdH} \\end{equation}  If time is discretised as $t=hk$, with $h$ a constant timestep, the sequence of increments  \\begin{equation} G_H (k) = B_H (k) - B_H(k-1) \\end{equation} is what is sometimes called fractional Gaussian noise, with the property that it is causally strongly connected with itself even at infinite temporal range. This is seen by calculating the expectation value of the correlation function of this signal in the limit of an infinite numbers of timesteps  \\begin{equation} \\lim _{k\\to \\infty} E\\left [ G_H(j)G_h(j+k)\\right] \\sim H(2H-1)k^{2H-2} \\neq 0 \\end{equation} for $H \\notin \\{0,1/2\\}.$ So stated more clearly, the number $H$ is an indication of the memory of the process: the fBM $B_H$ has increments that are not independent of each other.  The set $\\{ v_r, v_\\phi, B_r, B_\\phi \\}$ at each position and timestep is needed. The necessary MHD equations are written as stochastic differential equations, with $\\xi (t)$ as a source term.  \\subsection{Equations valid for general $B_0(r)=\\left \\{0,B_\\phi(r),B_z(r) \\right\\}$}  After linearising the perturbed set of MHD equations and selecting just the needed set of variables, it is obtained that  \\begin{equation} \\frac{\\partial \\rho _1}{\\partial t} + \\frac{\\rho _0}{r}\\frac{\\partial (rv_r)}{\\partial r} + v_r\\frac{d\\rho_0}{dr} = 0, \\end{equation}  \\begin{equation} \\frac{\\partial v_\\phi}{\\partial t} + v_r \\left (\\frac{u_k}{r}+\\frac{du_k}{dr}\\right) = \\frac{1}{\\mu _0 \\rho _0}\\frac{a_1}{r^2}\\frac{d(rB_\\phi)}{dr}, \\end{equation}  \\begin{equation} \\frac{\\partial b_\\phi}{\\partial t} = a_1 \\frac{d(u_k/r)}{dr}-\\frac{\\partial (v_r B_\\phi)}{\\partial r}, \\end{equation}  \\begin{equation} b_r = \\frac{a_1}{r}. \\end{equation}  The continuity equation provides $v_r$ in terms of $\\xi (t)$  \\begin{equation} \\rho _0 (r) \\frac{\\partial v_r (r,t)}{\\partial r} + \\left [ \\frac{\\rho _0 (r)}{r} + \\frac{d \\rho _0 (r)}{dr} \\right ] = - \\xi (t), \\end{equation} which when solved for the $r$ variation with $v_r(r_o,t)=0$ gives  \\begin{equation}\\label{eqS:vr} v_r (r,t) = \\frac{r^{7/8} (-r^2 + r_o^2) }{2 a_2} \\xi (t). \\end{equation}  \\subsection{Configuration I of initial magnetic field\\label{sect:config1}}  We study how a disk with a particular initial configuration will evolve when subjected to a stochastic density perturbation. We let $B_\\phi = 0$ and  \\begin{equation} B_z(r) = \\frac{b_0}{r}\\sin \\left [ \\frac{2\\pi (r-r_i)}{r_o-r_i} \\right ] \\end{equation}  The value of $b_0$ is calculated by fixing the plasma beta $\\beta = 2 p_0 \\mu_0 / B_0^2$ to a numerical value.  The $\\phi$ component for the equation of motion and induction simplify and become  \\begin{equation} \\frac{\\partial v_\\phi}{\\partial t} + v_r \\left (\\frac{u_k}{r}+\\frac{du_k}{dr}\\right) = 0, \\end{equation} and  \\begin{equation} \\frac{\\partial b_\\phi}{\\partial t} = a_1 \\frac{d(u_k/r)}{dr}. \\end{equation}  The solution for the magnetic field is, with $b_\\phi(r,t=0)=0$  \\begin{equation} b_\\phi (r,t) = -\\frac{3a_1}{2}\\frac{u_k}{r^2} t = - \\frac{3a_1}{2}\\Omega _0 t. \\end{equation}  For this case we find that the perturbed magnetic field is not a stochastic variable.  The equation for the azimuthal velocity becomes  \\begin{equation} \\frac{\\partial v_\\phi}{\\partial t} + \\frac{1}{4a_2}u_k r^{-1/8} (-r^2+r_o^2)A_\\rho B_H (t)=0 \\end{equation}  The stochastic quantities in this configuration are $v_r$ and $v_\\phi$ and have defining equations  \\begin{equation} v_r (r,t) = f_1 (r) B_H (t), \\end{equation}  \\begin{equation} \\frac{\\partial v_\\phi (r,t)}{\\partial t} = f_2(r) B_H(t), \\end{equation} with  \\begin{equation} f_1(r) = A_\\rho \\frac{r^{7/8} (-r^2 + r_o^2) }{2 a_2} \\end{equation} and  \\begin{equation} f_2(r) = -\\frac{A_\\rho}{4a_2}u_k r^{-1/8} (-r^2+r_o^2). \\end{equation}  We take advantage that there is no differential operator in $r$ and write the defining equations in update form for the temporal variation  \\begin{equation} v_\\phi (r,t+h) = v_\\phi (r,t) + f_2(r) \\sqrt{h} G_H (t) \\end{equation} with $v_\\phi (r,0)=0$.  The way that $a_1$ is assigned a numerical value in this case takes advantage of the fact that $b_\\phi(r,t=0)=0$, and that $b_r=a_1/r$ such that we start the simulation by imposing that at $t=0$, the perturbed quantities have the same plasma beta as the equilibrium.  \\subsection{Configuration II of initial magnetic field\\label{sect:config2}}  The shape of $B_z(r)$ is kept and $B_\\phi (r) = \\frac{b_0}{r}\\cos \\left [ \\frac{2\\pi (r-r_i)}{r_o-r_i} \\right ]$. $b_0$ is again calculated based on plasma beta considerations.  In this case the evolution equations for $v_\\phi$ and $b_\\phi$ are  \\begin{equation} v_\\phi (r, t+h) = v_\\phi (r,t) - f_3(r) h +f_4 (r)\\sqrt{h}G_H(t), \\end{equation}  \\begin{equation} b_\\phi (r,t+h) = b_\\phi (r,t) + f_5 (r) h - f_6 (r)\\sqrt{h}G_H(t)\\label{eq:bphi2}, \\end{equation} where  \\begin{equation} f_3(r)=-\\frac{1}{\\mu_0 \\rho_0(r)}\\frac{a_1}{r^2}\\frac{d(rB_\\phi(r))}{dr}, \\end{equation}  \\begin{equation} f_4(r)= - \\left ( \\frac{u_k}{r}+\\frac{du_k}{dr}\\right ) f_1(r), \\end{equation}  \\begin{equation} f_5(r) = a_1 \\frac{du_k/r}{dr}, \\end{equation}  \\begin{equation} f_6(r) = \\frac{dB_\\phi f_1}{dr}. \\end{equation}  With the assumptions of the model, it is clear from Eq.~\\ref{eq:bphi2} that if there is a non-zero initial azimuthal magnetic field, the perturbed magnetic field will be a stochastic process. ",
        "conclusions": "\\label{sect:Conclusions}  The appearance of a perturbed magnetized accretion disk was investigated for the realistic case of a stochastic density perturbation, the output was investigated as a function of the Hurst parameter of the input signal. Two possible configurations of an initial magnetic field were considered. The most important result is the set of values obtained for the spectral slope and the corresponding Bayesian probabilities ($\\alpha$ and $p_B$, shown in Tables~\\ref{tab:C1b100-5-20}-\\ref{tab:C2b1000-5-100}). The simulated curves are found to reproduce the PSD characteristics of observed IDV data.  Although left with important things to fine-tune, this toy model may turn out to be very useful in explaining IDV. Future research in this area would benefit from focusing on finding a physics-based analytical bijective relation between input perturbation and produced light curves."
    },
    "1402/1402.0867_arXiv.txt": {
        "abstract": "Star formation lies at the center of a web of processes that drive cosmic evolution: generation of radiant energy, synthesis of elements, formation of planets, and development of life. Decades of observations have yielded a variety of empirical rules about how it operates, but at present we have no comprehensive, quantitative theory. In this review I discuss the current state of the field of star formation, focusing on three central questions: what controls the rate at which gas in a galaxy converts to stars? What determines how those stars are clustered, and what fraction of the stellar population ends up in gravitationally-bound structures? What determines the stellar initial mass function, and does it vary with star-forming environment? I use these three question as a lens to introduce the basics of star formation, beginning with a review of the observational phenomenology and the basic physical processes. I then review the status of current theories that attempt to solve each of the three problems, pointing out links between them and opportunities for theoretical and numerical work that crosses the scale between them. I conclude with a discussion of prospects for theoretical progress in the coming years. ",
        "introduction": "\\label{sec:intro}  Star formation is one of the least understood phenomena in cosmic evolution. It is difficult to formulate a general theory for star formation in part because of the wide range of physical processes involved. The interstellar gas out of which stars form is a supersonically-turbulent, weakly-ionized plasma governed by non-ideal magnetohydrodynamics (MHD). This by itself would make star formation a difficult problem, since we have at best a partial understanding of subsonic hydrodynamic turbulence, let alone supersonic non-ideal MHD turbulence. The behavior of star-forming gas is obviously influenced by gravity, which adds complexity, and the dynamics of the interstellar medium (ISM) is also strongly affected by both continuum and line radiative processes. Finally, its behavior is influenced by a wide variety of chemical processes, including formation and destruction of molecules and dust grains (which changes the thermodynamic behavior of the gas) and changes in ionization state (which alter how strongly the gas couples to magnetic fields). As a result of these complexities, there is nothing like a generally agreed-upon theory of star formation as there is for stellar structure. Instead, we are forced to take a much more phenomenological approach.  Before diving into this phenomenology, however, it is worth pausing to consider the motivation for this review. Star formation has been the subject of a number of recent reviews, focusing on theory \\citep{mckee07a}, numerical simulations \\citep{klessen11a}, observations in the Milky Way and nearby galaxies \\citep{kennicutt12a}, and a number of other much more detailed topics \\citep[e.g.,][]{hennebelle12a, kruijssen14a, dobbs14a, offner14a, padoan14a, krumholz14a, tan14a}. Each of these reviews provides a valuable description of one or more aspects of the star formation process, and some aim at a much more comprehensive overview of the field. Replicating either the scope or the detail of this previous work is neither a useful exercise, nor, given the limitations of space and reader attention, a viable goal.  Part of the motivation for this review is simply to provide an update. While there are a number of very recent reviews of specialized topics within star formation, the last comprehensive review of star formation theory as whole, by \\citet{mckee07a}, is now six years old. This is a very long time in a fast-moving field like star formation. However, the aim of this review also differs from that of previous reviews in two ways.  First, in this review I provide a significantly more pedagogic introduction to the subject, particularly on the basic physics background that is often assumed or skipped in higher-level reviews. While there are several textbooks on star formation \\citep{stahler05a, ward-thompson11a, bodenheimer11a, schulz12a}, this material is not part of the standard graduate curriculum at most institutions, and there is little material available that bridges the gap between a textbook-level introduction and a specialist review. The discussion I provide here is intended to occupy that middle ground, and is aimed at readers who are not experts in either the observational phenomenology or background theory of star formation or the molecular ISM, but want a much shorter and higher-level introduction than is provided by a textbook. I divide this introduction into a review of the observational phenomenology (Section \\ref{sec:phenomenology}) and an introduction to the physics of the star-forming interstellar medium (Section \\ref{sec:theorybackground}). The latter in particular focuses on basic theory as much as recent results. Both these sections are intended to bring students and other non-experts up to speed, and may safely be skimmed or skipped by readers who already possess a deep familiarity with star formation.  The second goal of this review is to provide a (necessarily biased) perspective focused on what I consider the minimal elements required for a predictive theory of star formation, and to suggest a common approach to tackling them. This is an important difference from most of the more narrowly-focused reviews listed above, which strive to cover one aspect of star formation in great detail. My goal here is instead to step back and identify those questions that will need to be answered before star formation theory becomes like the theory of stellar evolution: a field that continues to hold unsolved problems, but one for which enough basic results have been established that researchers in other areas of astrophysics can make use of them with some confidence, and without the need for continual worry if even the zeroth-order results are robust. Reaching this point is necessary if we are ever to have confidence in extrapolating the results of star formation studies to galactic environments far-removed from those familiar to us from the nearby Universe.  Any predictive theory must provide a statistical description of the star formation process, and while many statistics can be defined, I focus on three in this review. First, given the large-scale properties of a galaxy or some subsection of it (e.g., gas and stellar distributions and kinematics, metal content), what will be the \\textit{star formation rate} (SFR) in that galaxy? Second, what will be the \\textit{spatial and temporal distribution} of the star formation, i.e.~how will the newborn stars be clustered together in space and time? Third, what will be the \\textit{mass distribution} of the resulting stars, the initial mass function (IMF)? Each of these questions could be (and has been) the basis of its own review, but as I argue below, they are inextricably linked, and must be solved together.  The remainder of this review is organized as follows. As mentioned above, I begin in Section \\ref{sec:phenomenology} with a review of some of the necessary observational background, and provide a similar introduction to the theoretical background in Section \\ref{sec:theorybackground}. I then review the three questions of the star formation rate, the clustering of star formation, and the mass distribution of stars in Sections \\ref{sec:sfr}, \\ref{sec:clustering}, and \\ref{sec:imf}, respectively, before attempting a rough synthesis and pointing the way for future work in Section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{From the SFR to Clustering to the IMF: Linking the Scales}  I hope that those who have made it this far through the review will have drawn a few conclusions of their own. One that should be obvious is that the big three problems of the SFR, stellar clustering, and the IMF are inextricably linked, and a successful theory to explain one of them is likely to either be a theory for all three, or at a minimum to require certain assumptions about all three. These linkages of the galactic to the stellar scale come about in several ways, and which of them are important depends on the theory. At a minimum, the formation of molecular clouds at large scales provides initial conditions for the subsequent formation of stellar clusters and fragmentation of gas into individual stars. The feedback produced by those stars then provides a boundary condition for the star formation on galactic scales. Realistically, the linkages go beyond that. For example, radiative feedback may well be important in setting the IMF, but its strength is affected both by the star formation rate, which determines the accretion luminosity, and by the spatial clustering of accreting stars, which determines how strongly the radiation from one star can influence its neighbors.  Trying to build theories that address all these scales at once has been somewhat at odds with the traditional approach to star formation theory, and even with many of the current models. For example, the ``top-down\" approach to building a theory for the star formation rate amounts to hypothesizing that everything important is determined by galactic scale processes, while everything on the cluster and smaller scale hidden behind one or two parameters. Conversely, models that predict the peak of the IMF in terms of a (turbulent) Jeans mass pre-suppose the presence of a cloud with a well-defined density or linewidth-size relation, and treat these as single numbers that are set at the galactic scale, and can be regarded as fixed parameters. While such simplifications can be powerful, one of the lessons that I hope readers take from this review is that we are starting to bump up against their limitations. Models that treat all the small-scale physics in the star formation rate as a single number seems to founder on how to handle the observed metallicity-dependence of the star formation rate, and models for the IMF that treat the galactic scale via a single density or sonic length have trouble explaining why the IMF does not vary more than we observe.  These and other examples suggest that future progress will depend on linking the scales by developing theories that are naturally multi-scale in nature. On the analytic side, the most promising venue for this appears to be through statistical theories of turbulence, which naturally embrace the multi-scale nature of the cold interstellar medium. Their power is illustrated by the fact that approaches of this sort have been featured prominently in the discussions above on each of the three big problems, sometimes even within the same theoretical framework. The major limitation of these models is that they are scale-free and at present have difficulty incorporating the necessarily messier physics of stellar feedback, which is decidedly non-scale free. We have seen, however, that turbulence alone is probably not sufficient to explain the star formation rate, explain the fraction of stars that form in bound clusters, or determine the position of the IMF peak. Some other ingredient is required, and stellar feedback is the natural suspect.  The numerical challenge is to build simulations that include enough physics to go persuasively from the galactic scale down to the point of resolving individual stars. Due to the huge dynamic range of the problem, a simulation that does this is likely to be able to run for only a very small fraction of a galactic dynamical time, but even such a short-duration run is likely to yield useful results. The real challenge will be building in the requisite physics. At present, galaxy-scale simulations tend to assume optically-thin heating and cooling and use either no chemistry or only highly-simplified chemical networks. They generally do not include radiative transfer to treat stellar radiation feedback, or indeed any other feedback processes except via a parameterized energy or momentum injection recipe. In contrast, small-scale simulations often include much more detailed treatments of stellar feedback that do not rely on recipes, but make a number of assumptions that are clearly inappropriate on the galactic scale -- for example that dust and gas are tightly coupled, or that the gas is fully molecular. There are also intermediate-scale simulations that treat the atomic to molecular transition, but generally do not include stellar feedback. At present, there are no simulation codes that include all the physical processes that are necessary to credibly simulate the ISM on scales from the galactic to the stellar. Building such a code must be a major goal of theoretical work on star formation in the next five years.  As an adjunct to this, there will also need to be significant work on connecting the output to observations. Some of the observational indicators discussed in this review are fairly straightforward, such as maps of the locations and numbers of protostars as compared to maps of dust mass and column density. Some, however, require significantly more interpretation, such as observations of optically thick molecular lines. Proper comparison with observations, and successful mapping from observed to physical quantities, will require post-processing simulation results with radiative transfer and chemistry codes.  \\subsection{Problems for the Next Few Years}  It seems appropriate to end this review with some suggestions for ``low-hanging fruit\": problems that, unlike the grand challenges we have just discussed, should be solvable in the next one to two years, and that would probably generate a large number of citations should an ambitious young postdoc or grad student come up with a solid answer. Most of these ideas have been mentioned before in the review, but hopefully summarizing them here will provide a quick cheat sheet of inspiration for projects.  \\begin{itemize} \\item Supernova feedback in a turbulent medium. The feedback-regulated model of star formation discussed in Section \\ref{sssec:feedbackmodel} has as one if its crucial unknowns the amount of momentum that is produced by expanding supernova blast waves. Present simulations for the most part do not reach the resolution required to determine this quantity in a reliable manner. Using an AMR code or similar, it should be possible to simulate star formation at a galactic scale and follow the collapse down to the point where supernova blast waves would be resolved, and then to set off supernovae (perhaps using a subgrid model for the star formation itself) and measure the amount of momentum they impart. Such a simulation would probably be too expensive to run for multiple galactic dynamical times, but just following it far enough for the supernova remnant to reach the momentum-conserving phase would provide a very useful estimate for a quantity whose current value is very poorly known. \\item Stellar clustering with gas removal. The turbulence models of Hennebelle and Chabrier and Hopkins make predictions for the spatial distribution of dense gas cores and young stars that appear at least roughly consistent with observations (Section \\ref{ssec:starstructure}), but say nothing about how gas is removed and how those stars evolve once feedback begins to remove gas. In contrast, Kruijssen presents a model of the response of a hierarchically-structured stellar population to gas removal, but, since his model relies only on the single-point density PDF, his treatment of the spatial arrangement of stars before gas removal is very crude, and he cannot easily make predictions for the cluster mass function (Section \\ref{sssec:clusterfrac}). It ought to be possible to put these elements together to come up with an improved theory that can simultaneously predict both the cluster mass function and the fraction of stars forming in clusters. \\item Isothermal collapse: points or filaments? As discussed in Section \\ref{sssec:imfpeakgal}, one major class of theories for the origin of the peak of the initial mass function is that it is set by turbulence. However, it has yet to be demonstrated that turbulence plus self-gravity produces a set of point masses with a definite mass scale. We know of at least one example of an isothermal fragmentation problem where the mode of collapse is to a singular filament rather than singular points, in which case there is no well-defined answer to the question of what is the characteristic mass of the fragments. Is this just a special, singular case, or does isothermal collapse generally produce networks of singular filaments rather than singular points? By the same token, is it possible to obtain a numerically-converged result for the gravitational fragmentation of an isothermal turbulent medium, or do such simulations invariably fail to converge and produce a mass spectrum that extends down to the resolution limit? \\item Localized radiative feedback in turbulent fragmentation models. As discussed in Section \\ref{sssec:imfpeaknoniso}, radiative feedback appears capable of modifying the way gas fragments, but this effect is not included in any of the broader theories of turbulent fragmentation. While the models can handle non-isothermal equations of state that apply everywhere, they cannot yet handle localized feedback. Extending them to do so, and checking how this affects their predictions for the IMF, would be very valuable. \\end{itemize}  Of course (as often happens) some of these problems may be much deeper and more difficult than they appear. As with this entire review, this list also reflects the biases of the author as to what is worth doing. Caveat emptor."
    },
    "1402/1402.3169_arXiv.txt": {
        "abstract": "{ Our understanding of the $s$-process nucleosynthesis in asymptotic giant branch (AGB) stars is incomplete. AGB models predict, for example, large overabundances of lead (Pb) compared to other $s$-process elements in metal-poor low-mass AGB stars. This is indeed observed in some extrinsically enhanced metal-poor stars, but not in all. An extensive study of intrinsically $s$-process enriched objects is essential for improving our knowledge of the AGB third dredge-up and associated $s$-process nucleosynthesis.} {We compare the spectral abundance analysis of the SMC post-AGB star J004441.04-732136.4 with state-of-the-art AGB model predictions with a main focus on Pb. The low S/N in the Pb line region made the result of our previous study inconclusive. We acquired additional data covering the region of the strongest Pb line.} {By carefully complementing re-reduced previous data, with newly acquired UVES optical spectra, we improve the S/N of the spectrum around the strongest Pb line. Therefore, an upper limit for the Pb abundance is estimated from a merged weighted mean spectrum using synthetic spectral modeling. We then compare the abundance results from the combined spectra to predictions of tailored AGB evolutionary models from two independent evolution codes. In addition, we determine upper limits for Pb abundances for three previously studied LMC post-AGB objects. } {Although theoretical predictions for J004441.04-732136.4 match the $s$-process distribution up to tungsten (W), the predicted very high Pb abundance is clearly not detected. The three additional LMC post-AGB stars show a similar lack of a very high Pb abundance.} { From our study, we conclude that none of these low-mass, low-metallicity post-AGB stars of the LMC and SMC are strong Pb producers. This conflicts with current theoretical predictions.} ",
        "introduction": "\\label{sect:intro}  During the asymptotic giant branch (AGB), stars undergo thermal pulses that may be followed by third dredge-ups during which freshly synthesized products are brought into the envelope, among them the carbon and $s$-process elements. Theoretical and observational evidence shows that the main neutron source in low-mass AGB stars (1-3 M$_{\\odot}$) is the $^{13}$C($\\alpha$,n)$^{16}$O reaction \\citep{straniero95,gallino98,abia02}. It is widely accepted that the $^{13}$C pocket originates in a region where protons from the H-rich convective envelope are mixed into the He-rich intershell after a thermal pulse. However, the physical mechanisms of both third dredge-up and partial mixing in the intershell remain poorly understood.\\newline The $s$-process nucleosynthesis is predicted to depend strongly on metallicity while the $^{12}$C component, hence the $^{13}$C neutron source, in the intershell are of primary origin. In metal-deficient environments ([Fe/H] $<$ -1), in proportion, more neutrons are available for each iron seed resulting in larger overabundances of heavy $s$-process elements with respect to lighter ones ([hs/ls]). However, to date no clear relation has been observed between [hs/ls] and metallicity (see e.g. Fig 10 in \\cite{vanaarle11}) in post-AGB stars. The end product of the $s$-process nucleosynthesis chain is the double magic lead $^{208}$Pb isotope, which is predicted to have large overabundances with respect to other s-elements in metal-poor conditions (see \\cite{gallino98},\\cite{goriely00} and \\cite[][and references therein]{lugaro12}). The detection of metal-deficient objects with strong Pb enhancement by e.g. \\cite{vaneck01,vaneck03} and \\cite{behara10} confirm these predictions. However, some low-metallicity extrinsically $s$-process enriched objects (i.e. binaries) were found without strong Pb overabundance \\citep{aoki01,vaneck03}. These are just a few illustrations that show that our understanding of the $s$-process nucleosynthesis is limited. Additional systematic observations are required to deduce $s$-process distributions, including Pb abundances, to improve our understanding of the third dredge-up and AGB nucleosynthesis mechanisms \\citep[e.g.][and references therein]{herwig05}.\\newline In \\cite{desmedt12} (paper I), we performed an extensive spectral abundance analysis of the Small Magellanic Cloud (SMC) post-AGB object \\object{J004441.04-732136.4} (hereafter abbreviated to J004441). As a 21 $\\mu$m source \\citep{volk11}, J004441 shows very strong $s$-process enhancement combined with a low metallicity of [Fe/H] = -1.34 $\\pm$ 0.32 and a modest C/O ratio of 1.9 $\\pm$ 0.7. A luminosity of 7600 $\\pm$ 200 L$_{\\odot}$ was derived with the known distance to the SMC. Comparison of the position of J00441 in the HR diagram to evolutionary post-AGB tracks of \\cite{vassiliadis94} resulted in an estimated initial mass of approximately 1.3 M$_{\\odot}$. Based upon the initial mass and metallicity, different theoretical models were calculated with two independent stellar evolution codes, namely the Mount-Stromlo Evolutionary code \\citep[][and references therein]{karakas10a} and the STAREVOL code \\citep[e.g.][and references therein]{siess07}. The theoretical predictions match the $s$-process distribution but fail to reproduce the high overabundances and the modest C/O ratio. All models predict strong Pb overabundances but unfortunately, the strongest spectral Pb line at 4057.807 $\\AA$ was only covered by a very low S/N spectrum, preventing an accurate Pb abundance determination with the data of Paper I. We therefore reanalysed the Pb abundance of J004441 using a newly obtained UVES spectrum which is described in this letter.\\newline In addition, we determine upper limits for the Pb abundances for three Large Magellanic Cloud (LMC) post-AGB objects. Accurate spectral abundance studies of these objects have been conducted by \\cite{vanaarle11}, but Pb was not included in their analysis. Comparison of the different Pb enhancements will provide insight into the intrinsic enrichment of metal-poor, low-mass AGB objects. \\newline In the following section, we report on the new data and Section~3 describes our analysis and abundance determination. We compare the results of J004441 with model predictions in Section~4, followed by analyses of the Pb abundances for the LMC objects in Section~5. We end with a brief discussion.\\newline ",
        "conclusions": "\\label{sect:disc} Using additional spectra of J004441, an $s$-process rich post-AGB star in the SMC, we obtained an upper limit of the Pb abundance. With [Pb/Fe] $<$\\,2.58,  Pb shows a similar or smaller overabundance than the $s$-process elements of the Ba-peak.  However, this upper limit is much smaller than predicted by our specific AGB models. The latter were tuned to the initial mass and metallicity of the object and focussed on a fit through the obtained overabundances of the Zr- and Ba-peak elements. This discrepancy is real and not due to the poor S/N of the spectrum. The $s$-process predictions of the tuned AGB models, based on successive thermal pulses with associated partial mixing and dredge-ups are not able to reproduce the low Pb abundance.\\newline The [Pb/Fe] results for the three LMC objects are similar to J004441: there is no predicted high Pb overabundance relative to the Ba-peak elements. Although having a strong $s$-process enrichment, the four low-mass, metal-deficient post-AGB stars of the Magellanic clouds all show the absence of the theoretically predicted strong Pb overabundances. This trend is similar to what is observed in some, but not all, extrinsically enriched objects in \\cite{aoki01} and \\cite{vaneck03}. Current  AGB models based on a $^{13}$C-pocket arising from diffusive overshooting at the base of the convective envelope during the third dredge-up have problems reproducing the observed low Pb abundances in metal poor stars, the spread in [hs/ls] at a given metallicity and/or the C/O ratios. This clearly indicates that some physical ingredients are missing in the description of this process. To improve the situation, we need to obtain full abundance patterns of many post-AGB stars in both the LMC and SMC, so that we can identify systematics in the observed patterns. Simultaneously, we need to explore alternative processes to explain the derived abundance profiles. In this study, we used the diffusive approach to form the $^{13}$C pocket, but the partial mixing of protons below the envelope may not be diffusive in essence but rather advective. For instance, when adopting a new mixing algorithm, \\cite{straniero06} were able to generate a more massive $^{13}$C pocket and consequently produce higher $s$-process surface enrichments in better agreement with observations. Rotation is also a key physical ingredient that may have a strong impact on the synthesis of $s$-elements. Depending on the degree of shear, the $s$-process nucleosynthesis can be completely inhibited (see e.g. \\cite{langer99}, \\cite{herwig03} and \\cite{siess04}) or partially activated \\citep{piersanti13}. Also the effect of internal gravity wave mixing \\citep{denissenkov03} on the $s$-process nucleosynthesis still needs to be explored. \\newline Additional work, both theoretically assisted by direct hydrodynamical simulations and observational with systematic abundance determinations of a whole population of post-AGB stars with well constrained distances and a spread in metallicity, is needed to better constrain the $s$-process mechanisms.  \\begin{figure}[tb!] \\resizebox{.97\\hsize}{!}{\\includegraphics{Fig3.png}} \\caption{Upper panel: Overview of [s/Fe] for the different objects. The values at Z=82 represent the derived Pb abundance upper limits. Lower panel: [s/Fe] for the different objects scaled to [La/Fe] at Z=57.}\\label{fig:comp} \\end{figure}"
    },
    "1402/1402.6927_arXiv.txt": {
        "abstract": "{The characterization of the optical properties of the atmosphere in the near UV, in particular the tropospheric aerosol stratification, clouds optical depth and spatial distribution are common in the field of atmospheric  physics, due to aerosol effect on climate, and also in cosmic rays physics, for a correct reconstruction of energy and longitudinal development of showers. The goal of the ARCADE project is the comparison of the aerosol attenuation measurements obtained with the typical techniques used in cosmic ray experiments (side-scattering measurement, elastic LIDAR and Raman LIDAR) in order to assess the systematic errors affecting each method providing simultaneous observations of the same air mass with different techniques. For this purpose we projected a LIDAR that is now under construction: it will use a 355 nm Nd:YAG laser and will collect the elastic and the N$_2$ Raman back-scattered light. For the side-scattering measurement we will use the Atmospheric Monitoring Telescope, a facility owned by the Colorado School of Mines and placed in Lamar (Colorado), the site where our experiment will take place.} ",
        "introduction": "The knowledge of the characteristics of atmospheric aerosols is an important topic in different scientific fields. It is closely related to the understanding of the complex mechanisms that determine the climatic variations of our planet, and nevertheless aerosols and clouds play a key role in the observations and study of High Energy Cosmic Rays ($E > 10^{12} eV$) and Ultra High Energy Cosmic Rays (UHECR) ($E > 10^{18} eV$) from both ground-based and space-based experiments \\cite{bib:atmo, bib:space}. \\\\ Common techniques adopted to observe extensive air showers (EAS) are based on the detection of the Cherenkov light \\cite{bib:CTA} and UV fluorescence light \\cite{bib:auger} that is emitted during the passage of secondary particles of the shower through the atmosphere. These techniques, that are sensitive to the development of cosmic ray showers in atmosphere, allow the identification of the primary particle that generated the shower and the reconstruction of its energy and mass composition from the study of the development of the light profile.\\\\ The fluorescence technique to detect EAS makes use of the atmosphere as a giant calorimeter whose properties must be continuously monitored to ensure a reliable energy estimate. Atmospheric parameters influence both the production of fluorescence light and its attenuation towards the detectors. The molecular and aerosol scattering processes that contribute to the overall attenuation of light in the atmosphere can be treated separately.  Molecular scattering is well described by the Rayleigh theory, and can be estimated by measuring atmospheric state variables. The scenario is more complex for aerosol attenuation, it depends on the particulate composition, size, and shape, and in general it can't be calculated analytically.  Aerosols are subject to significant variations on time scales of an hour and aerosol attenuation is the largest time dependent correction applied for air shower reconstruction. Due to the importance of this aspect of the atmospheric properties, EAS observatories need to develop ad-hoc atmospheric monitoring systems.  The ARCADE project aims to characterize the optical properties of the atmosphere in the near UV region and to develop predictive models through measurement campaigns in areas with different environmental characteristics. In detail, for the first time, different techniques of analysis used in Cosmic Rays experiments to study aerosol attenuation, cloud cover and cloud optical depth will be applied simultaneously on the same air mass to highlight the points of strength, limitations and specially reduce the systematic errors that affect each technique. We also plan to develop models that describe the stratification of aerosols in the areas of measurement and their dependence on temperature, wind, precipitation, humidity, in collaboration with climatology experts that will apply predictive models to these measurements. ",
        "conclusions": "ARCADE is a three years project, funded by the italian Ministero dell'Istruzione, dell'Universit\\`a e della Ricerca (MIUR), that started in 2012. After the completion of construction of the LIDAR, first measurement campaign will take place first in Turin (urban environment), then in L'Aquila (rural environment) at the end of 2013. Then the LIDAR will be shipped to Lamar (desertic environment) where the device will acquire data together with the existing facility AMT for a year. Results obtained by the project will be included in a database of free access to the scientific community."
    },
    "1402/1402.3505_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\hspace{1.1cm} Thanks to successful missions such as WMAP and Planck and their facility to discover and investigate cosmological effects with small amplitudes the interest to the problems of primordial recombination and radiation arising during this process increased significantly in the last decade. Many theoretical investigations have been performed to improve results (e.g. Grachev \\& Dubrovich 2008; Switzer \\& Hirata 2008; Hirata 2008) of earlier works (Zeldovich et al. 1968, Peebles 1968, Dubrovich 1975) and to predict new results (e.g. Grachev \\& Dubrovich 2011, Shaw \\& Chluba 2011). A few techniques have been suggested to observe primordial recombination radiation [PRR] (Dubrovich \\& Grachev 2004; Rubino-Martin et al. 2008; Sunyaev \\& Chluba 2008, 2009). These techniques are difficult to realize because of problems with cross-calibration of observation channels at different frequencies. On the other hand the technique of observation of small deviations from main spectrum at specific frequency but at the different angles is developed good enough from experimental point of view (e.g. space experiments Relikt-1, COBE, WMAP, Planck, Hershel) and continues to improve from theoretical point of view (see e.g. Chluba \\& Sunyaev 2008, Pitrou et al. 2010, Gruetjen \\& Shellard 2012, Shimon et al. 2012, Chluba et al. 2012, Challinor 2012 and references therein). Evident power of methods of CMB small angular fluctuation observations gives us a hope for further successful development of such methods of experimental cosmology to achieve sensitivities of about $10^{-8}$ (or even $10^{-10}$ in far future) of the background signal. One of these methods is the observation of the thermal Sunyaev-Zeldovich (SZ) effect. Classical SZ effect (Sunyaev \\& Zeldovich 1969) is the small ($10^{-5} - 10^{-3}$) change of CMB intensity (decrease in Rayleigh-Jeans range [$<218$ GHz] and increase in Wien range [$>218$ GHz]) due to comptonization of CMB during its transfer through hot (3 - 15 keV) intracluster medium (i.e. change of photon energies [increase mainly] due to inverse Compton scattering). In this paper we consider small corrections to the thermal SZ effect in galaxy clusters due to PRR and suggest a new method for observing PRR via measurements of these corrections by using anisotropy technique. A similar method allowing to measure the derivative of PRR spectrum using Doppler effect due to motion of observer relative to CMB rest frame will be described in the accompanying paper (Balashev et al. 2014 in preparation). If mentioned sensitivities will be achieved then such measurements will become an important source of information about PRR characteristics, cosmological parameters, and parameters of galaxy clusters. ",
        "conclusions": "The following quantities have been choosen as the characteristics of the main SZ effect and its corrections: \\\\1) Relative change of the occupation number (or intensity) $d = \\left(\\eta_{f}/\\eta_{i}-1\\right)$ (we will refer to this quantity simply as ``SZ effect''). \\\\2) Relative change of SZ effect due to distortion of background (CMB+CIB) by PRR $\\Delta d_{prr}/d_{bg}=\\left(d_{tot}/d_{bg}-1\\right)$, where $d_{tot}$ is the SZ effect for total spectrum (CMB+CIB+PRR), $d_{bg}$ is the SZ effect for spectrum of background (CMB+CIB). We will refer $d_{bg}$ as ``main SZ effect'' in case of general background and ``standard SZ effect'' for the range where CMB strongly dominates. \\\\3) Absolute difference of observed intensities $\\Delta I = \\left(I_{f}-I_{i}\\right)$, where $I_{i}$ is the intensity of radiation coming through empty vicinity near the cluster, $I_{f}$ is the intensity of radiation coming through intracluster medium. The quantity $\\Delta I$ is required to estimate perspectives of experimental registration of considered corrections from point of view of modern telescope sensitivity. Hereinafter we will refer to $\\Delta I$ as angular intensity difference.   \\subsection{Correction to SZ effect due to PRR in CMB range} The results for standard SZ effect $d_{bg}$ at different temperatures $T_e$ of electron gas in the galaxy cluster and at fixed Compton parameter $y_C=3\\cdot 10^{-4}$ are presented in the bottom panel of Fig. \\ref{fig1} (it should be noted that curves corresponding to different temperatures of electron gas are overlapped). The results for relative change $\\Delta d_{prr}/d_{bg}$ of SZ effect at same $T_e$ and $y_C$ values are presented in the top panel of Fig. \\ref{fig1}. Dependence of $\\Delta d_{prr}/d_{bg}$ on observation frequency $\\nu$ shows quasi-oscillations. These oscillations are originated from two phenomena: \\\\1) Lines in the initial spectrum (i.e. spectrum of incident PRR). \\\\2) Dependence of SZ effect on derivatives (high-orders also) of the initial spectrum. It can be shown by means of Taylor series of initial occupation number $\\eta_{i}$ in expression (\\ref{Redistribution_relation}). More evidently and qualitatively this dependence can be shown from the differential equation of Kompaneets approximation (e.g. Sunyaev \\& Zeldovich 1969): \\begin{equation} {\\partial \\eta \\over \\partial y_C}={x_{e}^{-2}}{\\partial \\over \\partial x_{e}} x_{e}^4 {\\partial \\eta \\over \\partial x_{e}} \\label{Kompaneets_equation} \\end{equation} where $x_{e}=h\\nu/\\left(k_BT_e\\right)$ is dimensionless frequency. Note that the kernel (\\ref{Kompaneets_kernel}) is the solution of equation (\\ref{Kompaneets_equation}). Let us consider the simple example of the initial spectrum: power spectrum $\\eta_{i}\\sim x_{e}^{-\\alpha}$. Then SZ effect can be estimated by the following formula: $d=\\alpha\\left(\\alpha-3\\right)y_C$ (e. g. Sunyaev \\& Zeldovich 1980). This formula can also be used in the case of slow-changing spectral index $\\alpha$  (when $\\partial \\alpha / \\partial \\ln x_e \\ll \\alpha \\left(\\alpha-3\\right)$). Thus alternating-sign behaviour of value $\\Delta d_{prr}/d_{bg}$ can easily be demonstrated by the following relation: \\begin{equation} \\left({\\Delta d_{prr} \\over d_{bg}}\\right)_{K} \\simeq {\\eta_{prr}\\over \\eta_{bg}}{\\alpha_{prr}\\left(\\alpha_{prr}-3\\right) \\over \\alpha_{bg}\\left(\\alpha_{bg}-3\\right)} \\label{D_d_d_analit} \\end{equation} where subscript ``K'' denotes ``Kompaneets approximation'', $\\eta_{bg}$ is the occupation number of background radiation, $\\eta_{prr}$ is the occupation number of PRR spectrum (i.e. the intitial occupation number is $\\eta_{i}=\\eta_{bg}+\\eta_{prr}$), $\\alpha_{bg}\\simeq -\\partial \\ln \\eta_{bg} / \\partial \\ln x_e$, $\\alpha_{prr}\\simeq -\\partial \\ln \\eta_{prr} / \\partial \\ln x_e$. One can see from equation (\\ref{D_d_d_analit}) that $\\Delta d_{prr} / d_{bg}$ changes its sign at the points where $\\alpha_{prr}\\simeq 0$ or 3 and $\\alpha_{bg}\\simeq 0$ or 3.  Amplitude of these quasi-oscillations ranges from $2\\cdot 10^{-9}$ (at $\\nu=100$ GHz) to $5\\cdot 10^{-7}$ (at $\\nu=0.3$ GHz). Quasi-period $\\Delta \\nu /\\nu$ of quasi-oscillations (i.e. period containing full oscillation) changes from $\\simeq 0.1$ (at $\\nu=0.3$ GHz) to $\\simeq 0.5$ (at $\\nu=100$ GHz). The features at $\\nu \\simeq 218$ GHz is connected only with the fact that SZ effect $d_{bg}$ goes to zero at this frequency (see bottom panel of Fig. \\ref{fig1}). Fig. \\ref{fig1} shows that value $\\Delta d_{prr}/d_{bg}$ depends on temperature of electron gas in the galaxy cluster: the hotter the gas, the less the magnitude of $\\Delta d_{prr}/d_{bg}$. This is the result of the kernel shape. One can see that value $\\Delta d_{prr}/d_{bg}$ as a function of frequency $\\nu$ has unique shape depending on parameters of PRR and temperature of electron gas. This should allow researchers to identify reliably observational manifestations of the effect of thermal corrections to SZ due to PRR.  The results of calculations of ``standard'' SZ effect $d_{bg}$ at different values of Compton parameter $y_C$ of the galaxy cluster and at fixed temperature $T_e$ of electron gas are presented in the bottom panel of Fig. \\ref{fig2}. The results of calculations of relative change $\\Delta d_{prr}/d_{bg}$ of SZ effect at the same values of $y_C$ and $T_e$ are presented in the top panel of Fig. \\ref{fig2} (it should be noted that curves corresponding to different Compton parameters are overlapped). Also in the top panel of Fig. \\ref{fig2} the relative amplitude of CMB distortions due to PRR is presented.  Here one can emphasize two interesting facts: 1) ``Amplification'' of distortion; 2) Independence of $\\Delta d_{prr}/d_{bg}$ on Compton optical depth.  First of them is the following: If we consider CMB intensity $I_{bg}$ as noise (background) $N$, PRR intensity $I_{prr}$ as useful signal $S$, galaxy cluster as natural filter, SZ effect $d_{bg}$ of CMB as noise $N'$ resulting from pass of noise $N$ through the filter, and correction $\\Delta d_{prr}$ to SZ effect due to PRR as useful signal $S'$ resulting from pass of signal $S$ through the filter, then we would observe an increase (or decrease) of signal-to-noise ratio from $S/N=I_{prr}/I_{bg}=\\eta_{prr}/\\eta_{bg}$ to $S'/N' = \\Delta d_{prr}/d_{bg}$. In the Kompaneets approximation with slow-changing spectral indecies $\\alpha$ corresponding amplification coefficient $C_{K}$ (subscript ``K'' denotes ``Kompaneets'') can be estimated by the following relation resulting from equation (\\ref{D_d_d_analit}): \\begin{equation} C_{K}={S'/N' \\over S/N}\\simeq{\\alpha_{prr}\\left(\\alpha_{prr}-3\\right) \\over \\alpha_{bg}\\left(\\alpha_{bg}-3\\right)} \\label{C_K} \\end{equation} One can see from equation (\\ref{C_K}) that coefficient of ``amplification'' depends on sharpness of derivative of initial correction spectrum. For example in the case of correction in Rayleigh-Jeans part of CMB one has $\\eta_{bg}\\sim {x_e^{-1}}$, and $\\alpha_{bg}=1$. Therefore ``amplification'' coefficient can be rougly estimated as $|C_{K}|\\simeq\\left|\\alpha_{prr}\\left(\\alpha_{prr}-3\\right)/2\\right|$. For $\\alpha_{prr}$ achieving values about 8, $|C_{K}|$ achieves values of about 20 (see top panel of Fig. \\ref{fig2}). Exact estimate of the ``amplification'' coefficient $C_{amp}$ by using generilized approach gives us values (for amplitudes, i.e. at the points of maxima of $|\\Delta d_{prr}/ d_{bg}|$) of about $0.25$ - $10$ in Rayleigh-Jeans part of CMB depending on frequency (see top panel of Fig. \\ref{fig2}).  Second fact is the following: $\\Delta d_{prr}/d_{bg}$ does not depend on Compton optical depth $\\tau$ considerably (in the main order). This is clearly demonstrated by the overlap of curves $\\Delta d_{prr}/d_{bg}$ (see top panel of Fig. \\ref{fig2}) corresponding to different values of $y_{C}$ at fixed $T_{e}$. It can be shown by using formulae (\\ref{Redistribution_relation}) and (\\ref{general_kernel}) that the following relation is valid: \\begin{equation} {\\Delta d_{prr} \\over d_{bg}}= \\left({\\eta_{bg}\\over \\eta_{bg}+\\eta_{prr}} {F\\left(T_{e},\\left[\\eta_{bg}+\\eta_{prr}\\right]\\right) \\over F\\left(T_{e},\\eta_{bg}\\right)}-1\\right)+O\\left(\\tau\\right) \\label{Delta_d_d_no_tau} \\end{equation} where the second term (depending on $\\tau$) is much less than first term (independing on $\\tau$), function $F$ is given by the following formula: \\begin{equation} F\\left(T_{e},\\eta_{i}\\right)=\\int_{-\\infty}^{+\\infty}\\eta_{i}\\left(\\nu\\exp(-s)\\right)P_{1}\\left(-s, T_{e}\\right)ds \\end{equation} Formula (\\ref{Delta_d_d_no_tau}) shows that the knowledge of $\\Delta d_{prr}/d_{bg}$ gives us an opportunity to determine the temperature $T_{e}$ of electron gas in the galaxy cluster independently of Compton optical depth value $\\tau$. In other words formula (\\ref{Delta_d_d_no_tau}) shows how to exclude the dependence of experimantal data (obtained by means of observations of SZ effect from radiation with non-typical spectra) on Compton optical depth $\\tau$ in the main order. It should be noted that the determination of relative corrections $\\Delta d_{prr}/d_{bg}$ is more complicated than determination of SZ effect $d=\\Delta I / I_{i}$ (where both $\\Delta I$ and $I_{i}$ are the observable quantities), because during experiment one can observe only total intensity of radiation (or occupation number $\\left(\\eta_{bg}+\\eta_{prr}\\right)$) and correspondingly can calculate only total SZ effect $\\left(d_{bg}+\\Delta d_{prr}\\right)$, but not ``standard'' SZ effect $d_{bg}$ of background radiation $\\eta_{bg}$ (this situation is completely similar to one with possible measurements of PRR lines distorting CMB spectrum). Thus determination of $\\Delta d_{prr} / d_{bg}$  demands special treatment to separate useful signal $\\Delta d_{prr}$ and noise signal $d_{bg}$ from the total signal $d_{tot}$. However, this treatment can depend on type of the initial spectrum, features of expected results, quality of data and others. For example, in the case of determination of $\\Delta d_{prr}$ due to PRR, one can average the total SZ effect $d_{tot}$ over frequency on the scales larger than quasi-period of quasi-oscillations (see Fig. \\ref{fig1} and \\ref{fig2}) and associate this value with $d_{bg}$. Also additional assumptions (e.g. assumption on the exact knowledge of CMB spectrum [temperature $T_0$] and its SZ effect) can help us to make a separation of useful and noise signals.  The results of calculations of intensities $I_{i}$ and their angular differences $\\Delta I$ for different initial spectra (background [CMB mainly] and PRR) at different values of cluster electron gas temperature $T_e$ and at fixed Compton parameter $y_C=3\\cdot10^{-4}$ are presented in Fig. \\ref{fig3}. Angular intensity differences for PRR also show marked quasi-oscillations (as well as $\\Delta d_{prr}/d_{bg}$). Amplitude of these quasi-oscillations has values in the range from $6\\cdot10^{-29}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $6$ $\\mu$Jy/ster at $\\nu=0.3$ GHz) to $6.7\\cdot 10^{-26}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $6.7$ mJy/ster at $\\nu \\simeq 337$ GHz). Fig. \\ref{fig3} also shows that in considered frequency range the larger angular intensity differences for PRR correspond to the lower electron gas temperature while the angular intensity differences for CMB are almost independent of the temperature of electron gas in the galaxy cluster (at fixed Compton parameter $y_C$).  The results of calculations of intensities $I_{i}$ and their angular differences $\\Delta I$ for different initial spectra (background [CMB mainly] and PRR) at different values of Compton parameter $y_C$ and at fixed cluster electron gas temperature $T_e=5$ keV are presented in Fig. \\ref{fig4}. One can see that at fixed cluster electron gas temperature $T_e$ angular intensity differences are almost directly proportional to $y_C$, i.e. directly proportional to Compton optical depth $\\tau$ (that is expected for small values of $\\tau\\ll 1$). For the value of Compton parameter $y_C=10^{-3}$ (at $T_e=5$ keV this corresponds to $\\tau\\simeq 0.1$) the amplitude of quasi-oscillations of PRR angular intensity difference achieves about $2.2\\cdot 10^{-25}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $22$ mJy/ster at $\\nu \\simeq 337$ GHz).   \\subsection{Correction to SZ effect due to PRR in CMB-CIB range} Calculation results for different characteristics of SZ effect in the frequency range $\\nu = 700 - 5000$ GHz are presented in Figs \\ref{fig5} - \\ref{fig9}.  It should be noted that CIB spectrum is not well known. We have used the data by Kneiske \\& Dole (2008) but the real spectrum of CIB can differ from used one significantly (it has strong cosmological anisotropy and is liable to the influence of local point sources). Thus the following results on relative change $\\Delta d_{prr}/d_{bg}$ of main SZ effect (see Figs \\ref{fig5}, \\ref{fig6}) serve only to estimate an order of magnitude of considered effect and to find the main features of behaviour of $\\Delta d_{prr}/d_{bg}$. For precise estimate of $\\Delta d_{prr}/d_{bg}$ a new calculation taking specific conditions (especially the intensity of CIB in specific direction for choosen object) into account should be carried out. For calculations of main SZ effect and correction to it (Figs. \\ref{fig5} - \\ref{fig8}) value of redshift of galaxy cluster $z$ has been choosen equal to 0.2.  The results of calculations of SZ effect $d_{bg}$ for background radiation (CMB+CIB) at different temperatures $T_e$ of electron gas in galaxy cluster and at fixed Compton parameter $y_C=3\\cdot 10^{-4}$ are presented in the bottom panel of Fig. \\ref{fig5}. The results of calculations of relative change $\\Delta d_{prr}/d_{bg}$ of SZ effect at the same values of $T_e$ and $y_C$ are presented in the top panel of Fig. \\ref{fig5}.  The relative change $\\Delta d_{prr}/d_{bg}$ has typical values $10^{-8} - 10^{-5}$ except narrow frequency bands where the main SZ effect $d_{bg}$ approaches zero. In these bands the value $\\Delta d_{prr}/d_{bg}$ becomes large (infinite in the points of zeros of main SZ effect) but it is only consequence of definition of $\\Delta d_{prr}/d_{bg}$ and does not have any physical sense. Nevertheless in some bands where main SZ effect has notable values ($d_{bg}\\le 10^{-4}$) the relative correction $\\Delta d_{prr} / d_{bg}$ can reach values larger than $10^{-5}$. From Fig. \\ref{fig5} one can see that in the difference with the behaviour of correction $\\Delta d_{prr}/d_{bg}$  in the range 0.3 - 700 GHz its behaviour in the range 700 - 5000 GHz depends on gas temperature $T_e$ significantly. The main reason of such dependence is the more complex (from SZ effect point of view) behaviour of background (CMB+CIB) spectrum in the range 700 - 5000 GHz in comparison with one in the range 0.3 - 700 GHz (almost pure CMB). This leads to appreciable dependence of main SZ effect on temperature $T_e$ (see bottom panel of Fig. \\ref{fig5}). For example frequency of first zero in the range 700 - 5000 GHz depends on $T_e$ approximately as $~T_{e}^{0.2}$. That in turn leads to the significant dependence of frequency values of poles (and bands of rapid growth) of correction $\\Delta d_{prr}/d_{bg}$ on temperature $T_{e}$. Such dependencies of main SZ effect and its correction on electron gas temperature $T_e$ may be one more useful instruments for determination of $T_e$ and CIB spectrum in the future experiments.  The results of calculations of SZ effect $d_{bg}$ for background radiation (CMB+CIB) at different values of Compton parameter $y_C$ and at fixed temperatures $T_e=5$ keV of electron gas in the galaxy cluster are presented in the bottom panel of Fig. \\ref{fig6}. The results of calculations of relative change $\\Delta d_{prr}/d_{bg}$ of SZ effect at the same values of $T_e$ and $y_C$ are presented in the top panel of Fig. \\ref{fig6}.  As well as for the correction to the SZ effect $\\Delta d_{prr}/d_{bg}$ in the range 0.3 - 700 GHz one can see from Fig. \\ref{fig6} that at fixed $T_{e}$ the main SZ effect $d_{bg}$ depends on Compton parameter $y_{C}$ linearly (i.e. depends on Compton optical depth $\\tau$ linearly) while correction $\\Delta d_{prr}/d_{bg}$ does not depend on Compton parameter $y_{C}$ in the main order (i.e. does not depend on Compton optical depth $\\tau$ in the main order). In the range 700 - 1100 GHz one can see that $\\Delta d_{prr} / d_{bg}$ is practically equal to relative distortion of background spectrum due to PRR. It is easy to show analitically that in the first order $\\Delta d_{prr} / d_{bg} \\simeq \\left(\\Delta I_{prr}/\\Delta I_{bg} - I_{prr}/I_{bg}\\right)$. Thus this effect takes place in the case of weak ``amplification'', i.e. when $C_{amp}\\ll 1$ (when disturbing spectrum $I_{prr}$ is smooth in comparison with background spectrum $I_{bg}$) and correspondingly  $\\Delta I_{prr}/\\Delta I_{bg} \\ll I_{prr}/I_{bg}$ and one can write $\\Delta d_{prr} / d_{bg} \\simeq  - I_{prr}/I_{bg}$.  The results of calculations of intensities $I_{i}$ and their angular differences $\\Delta I$ for the different initial spectra (background [CMB + CIB] and PRR) at different values of cluster electron gas temperature $T_e$ and at fixed Compton parameter $y_C=3\\cdot10^{-4}$ are presented in Fig. \\ref{fig7}. As well as in CMB range the angular intensity differences for PRR show marked irregular quasi-oscillations. Amplitudes of these PRR intensity differences have values in the range from $6.2\\cdot10^{-27}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $0.62$ mJy/ster at $\\nu \\simeq 773$ GHz) to $7.6\\cdot 10^{-26}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $7.6$ mJy/ster at $\\nu \\simeq 1700$ GHz). Unlike PRR intensity difference in the range 0.3 - 700 GHz (see Fig. \\ref{fig3}) the Fig. \\ref{fig7} shows that in the range 700 - 5000 GHz PRR intensity difference has more complex non-monotonic (in different frequency bands) dependence on electron gas temperature $T_e$. The background intensity differences also show appreciable dependence on temperature of electron gas $T_e$ (at fixed Compton parameter $y_C$).  The results of calculations of intensities $I_{i}$ and their angular differences $\\Delta I$ for the different initial spectra (background [CMB + CIB] and PRR) at different values of Compton parameter $y_C$ and at fixed cluster electron gas temperature $T_e=5$ keV are presented in Fig. \\ref{fig8}. As well as in the range 0.3 - 700 GHz Fig. \\ref{fig8} shows that at fixed cluster electron gas temperature $T_e$ angular intensity differences are almost directly proportional to $y_C$, i.e. directly proportional to Compton optical depth $\\tau$ in the main order. For the value of Compton parameter $y_C=10^{-3}$ (at $T_e=5$ keV this corresponds to $\\tau\\simeq 0.1$) the amplitude of quasi-oscillations of PRR angular intensity difference achieves about $2.5\\cdot 10^{-25}$ erg$\\cdot$cm$^{-2}$s$^{-1}$ster$^{-1}$Hz$^{-1}$ (i.e. $25$ mJy/ster at $\\nu \\simeq 1700$ GHz).  The last picture Fig. \\ref{fig9} illustrates behaviour of main SZ effect $d_{bg}$ (the bottom panel) and correction $\\Delta d_{prr} / d_{bg}$ at different values of redshifts of the galaxy cluster. One can see from Fig. \\ref{fig9} that at frequencies $\\nu$ larger than 1700 GHz both main SZ effect $d_{bg}$ and correction to it $\\Delta d_{prr} / d_{bg}$ show significant dependence on redshift of the galaxy cluster. Such behaviour differs from the common SZ effect in CMB range. The reason of this difference is non-trivial dependence of CIB spectrum on redshift $z$ due to radiation sources emitting in CIB range during evolution of the Universe, while CMB evolution is described by the simple relation of frequency redshifting $\\nu = \\nu_0 (1+z)$. This gives us possibility to determine the redshift of observed galaxy cluster, to separate its proper radiation in considered range from cosmological background (see also Rephaeli 1995) and to determine contribution to background from IR-sources in the late epochs ($<z$). Joint analysis of SZ effect and X-ray observations (independent determination of temperature $T_e$ of electron gas) will allow researchers to find essential limits of these characteristics.  \\vspace{0.3cm} \\hspace{0.0cm} {\\bf Acknowledgements} Authors are grateful to R. A. Sunyaev for careful reading and comments on this manuscript. This investigation has been performed with partial support of RFBR grant 13-02-12017-ofi-m and grant \"Leading Scientific Schools of Russia\" NSh-294.2014.2. SAB thanks RF Presedent Programme (grant MK-4861.2013.2)."
    },
    "1402/1402.4113_arXiv.txt": {
        "abstract": "We consider the conditions within a Gamma-Ray Burst (GRB) emission region that is Poynting flux dominated. Due to the enormous magnetic energy density, relativistic electrons will cool in such a region extremely rapidly via synchrotron. As there is no known mechanism that can compete with synchrotron it must be the source of the prompt sub-MeV emission. This sets strong limits on the size and Lorentz factor of the outflow. Furthermore, synchrotron cooling is too efficient. It overproduces optical and X-ray as compared with the observations. This overproduction of low energy emission can be avoided if the electrons are re-accelerated many times ($\\gtrsim 5\\times 10^4$) during each pulse (or are continuously heated) or if they escape the emitting region before cooling down. We explore the limitations of both models practically ruling out the later and demonstrating that the former requires two different acceleration mechanisms as well as an extremely large  magnetic energy to Baryonic energy ratio. To be viable, any GRB model based on an emission region that is Poynting flux dominated must demonstrate how these conditions are met. We conclude that if GRB jets are launched magnetically dominated they must dissipate somehow most of their magnetic energy before they reach the emission region. ",
        "introduction": "\\label{Int} It is generally accepted that Gamma Ray Bursts (GRBs) are powered by ultra relativistic jets ( with a Lorentz factor $\\Gamma\\geqslant100$) carrying large isotropic equivalent luminosities ($L_{iso}\\approx10^{53}$erg/sec) which are dissipated at large distances ($r\\approx10^{13}-10^{17}$cm) from the central engine. Whether the composition of the jets is: Baryonic  \\citep{Shemi(1990)} or magnetically (Poynting flux) dominated \\citep{Usov(1992),Thompson(1994),Meszaros(1997),Lyutikov(2003)} is one of the major puzzles concerning GRBs.   The thermal pressure at the base of an AGN jet is insufficient to support a bayronic outflow and therefore,  blazars' jets must be magnetically dominated. However, the strong IC component observed in blazers shows that this magnetic energy must have dissipated efficiently before the emission regions and the latter are baryonic dominated \\citep[see e.g.][]{GhiselliniTavecchio09}. The situation for GRBs could be very different as with smaller size engines the thermal pressure in the base could be sufficient to drive baryonic outflows. However, modeling of some GRBs' engines based on accretion disks suggest that the Poynting flux jet power is much stronger than the thermal driven outflow that derived from neutrino annihilation \\citep[see e.g. ][]{Kawanaka(2013)}. Furthermore, the fact that blazars jets are magnetically driven, suggests that this might also be the case for GRBs. On the other hand, the fact that unlike blazers GRBs don't have a strong high energy IC component suggests strong magnetic fields within their emission region.   The most direct clues on the nature of the jets arise from the observed prompt sub-MeV spectrum. As such, deciphering the nature of GRB jets is closely related to another major puzzle concerning GRBs - what is the radiation mechanism that produces the prompt emission. The situation is somewhat inconclusive. A strong upper limit on the thermal component  \\citep[e.g. GRB 080916C][]{Zhang(2009)} suggests a Poynting flux dominated outflow. On the other hand possible detection of $5-10\\%$ of the total energy in a thermal component \\citep[e.g. GRBs: 100724B, 110721A, 120323A][]{RydePeer(2009),Ryde(2010),Guiriec(2011),Guiriec(2013),Axelsson(2012)} which is not expected in a Poynting flux outflow, indicate a baryonic outflow.   The overall non thermal spectrum has led to the suggestion that the prompt emission is produced by synchrotron mechanism \\citep{Katz(1994),Rees(1994),Sari(1996),Sari(1998)}. This was supported by the reasonable association of the afterglow emission with synchrotron radiation. However, observations of many bursts showing lower energy spectral slopes steeper than the ``synchrotron line of death\" \\citep{Crider(1997), Preece(1998), Preece(2002)} have led to the suggestion that the prompt emission must be produced by another mechanism.   However, a refined time dependent spectral analysis in several bright bursts \\citep{Guiriec(2012)} has suggested that the situation might be much more complicated and GRB spectra may be better fitted with a multi component model instead of the classical ``Band function\". In the multi component model, the spectrum is composed of a superposition of a Band function peaking at few hundred keV and carrying most of the energy, a black body component at a few tens of keV with $5-10\\%$ of the energy and in some cases an additional power law component extending up to a few hundred MeV and carrying up to $40\\%$ of the total energy. In these fits, the lower energy spectral slope of the Band function becomes softer compared with the Band only fits, and is consistent with slow cooling synchrotron (i.e. the synchrotron ``line of death\" problem is removed). The direct detection of a weak photospheric component in those GRBs, suggests that the ``thermal signature\" is sub-dominant and it is not the main radiation source in operation during the prompt. These later results support synchrotron emission as the main radiation mechanism while requiring additional mechanisms to produce the other components.   In this paper, we explore these two strongly connected issues: the composition of GRB jets and the prompt radiation mechanism. Clearly, a Poynting flux dominated jet must dissipate a significant fraction of its magnetic energy and produce relativistic electrons that would emit efficiently the radiation at the emission zone. We consider, first, a simple one zone model for the Poyinting flux jet. Initially (the moment dissipation starts), $\\epsilon_B\\sim1$, in such scenario. After a fraction $\\epsilon_{diss}$ of the magnetic energy is dissipated, $\\epsilon_{B, after}=1-\\epsilon_{diss}$ remains as a magnetic energy. A fraction of the dissipated energy goes into electrons and a fraction of that is radiated away: $\\epsilon_{rad} <\\epsilon_e<\\epsilon_{diss}$.  \\citet{Panaitescu(2002), Granot(2006)} and \\citet{Fan(2006)} show that the radiative efficiency of GRBs is of order unity:  $\\epsilon_{rad}\\sim 0.2-0.8$. Therefore $\\epsilon_{diss}\\gtrsim 0.2-0.8$ and the ratio of magnetic to electron energy density at the emitting region is $\\epsilon_{B,after}/\\epsilon_e<(1-\\epsilon_{diss})/\\epsilon_{diss} \\sim 0.25-4$. Depending on the time it takes for the magnetic energy dissipation to take place there are two possibilities. If the dissipation is very rapid (compared to synchrotron cooling), then the actual radiation process takes place in a low magnetic field region and the radiation process ``is unaware\" of the original Poynting flux dominated nature of the outflow. This is similar to the situation in AGNs, where the magnetic energy within the emission region is sub-dominant. On the other hand, if the dissipation is slower, then the electrons effectively emit most of their energy in a region where the magnetic fields are still relatively high. We focus on this latter possibility in this paper, and we find constraints that should be imposed on such a scenario. These constraints have to be satisfied by any radiation process within a Poynting flux dominated emission region in order for it to be viable.   As relativistic electrons are essential to produce the observed prompt emission (regardless of the specific emission mechanism) then due to the high magnetic energy density, a significant synchrotron signal (not necessarily in the sub-MeV band) must be produced as well. We compare the expected synchrotron fluxes with upper limits on prompt observations at the optical, X-ray and GeV bands, and show that in order for the synchrotron flux to be consistent with observations, the gamma-ray process should be unrealistically efficient. Thus, in a magnetically dominated emission region, synchrotron emission must be the main emission mechanism during the prompt phase. This sets numerous limits on the conditions, such as the radius and Lorentz factor of the emitting regions.  Additionally, as the synchrotron is extremely efficient, the electrons will cool rapidly and in addition to producing a $\\nu^{-1/2}$ low energy spectral slope they will over produce optical and X-ray. This is inconsistent with observational limits within these energy bands. Thus, this low energy fast cooling tail must be avoided. In the last sections of this work we consider a few ways to avoid this problem. For example  it has been suggested that this inconsistency may be alleviated if electrons are re-accelerated many times during a single pulse however this may lead, in turn, to other problems. We also consider three simple two-zone models in which the magnetic field is inhomogenous. These toy-models sereve as demonstrating the intrinsic limitations of such configurations and highlight issues that must be addressed by any model of this kind.    This paper is organized as follows. In \\S \\ref{GenApp} we discuss the basic concepts and the parameter phase space of the model. In \\S \\ref{synsig} we discuss the general characteristics of synchrotron emission, focusing on the typical frequencies and the cooling timescale. In \\S \\ref{generalconstraints}, we find the constraints on the parameter space from observed limits on the flux in different bands. We discuss the implications of the results on the synchrotron emission mechanism, including several inhomgenous toy models in \\S \\ref{synmodel} and we conclude and summarize in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} Magnetically dominated jets need to dissipate some fraction of their energy to produce the observed radiation. However, it is unclear where the dissipation takes place and how efficient is it. In order to explore this issue we have explored the conditions within magnetically dominated emission regions, assuming that all the dissipation takes place right there. We note that even if there is significant dissipation before the emitting region, the limits on the cooling time we find in the paper (discussed below) may be viewed as limits on the allowed dissipation rate. Slower dissipation rates will effectively result in a magnetically dominated emission region, subject to the constraints found in this paper. Alternatively the dissipation should not lead directly to particle acceleration and this will provide a ``buffer\" during which the magnetic field decreases but rapid electrons are not available to emit synchrotron.   We considered a general emission mechanism, assuming only that relativistic electrons are involved and focusing on the inevitable synchrotron signature of these relativistic electrons in the presence of the strong magnetic fields. We find that for $\\Gamma\\lesssim600$ relativistic electrons that are accelerated on a time scale shorter than the dynamical time scale, will cool rapidly by synchrotron emission. Therefore, any radiation mechanism proposed to explain the sub-MeV peak that relies on relativistic electrons should be faster than synchrotron. As such a mechanism is not known, synchrotron is the main emission mechanism in magnetically dominated systems. If synchrotron produces the prompt sub-MeV emission, the electrons cool extremely rapidly. For example, within an outflow with a typical Lorentz factor $\\G=300$, the cooling time associated with the gamma ray emitting process, is at least one order of magnitude shorter than the pulse time scale. The typical cooling frequency will be below the X-rays and in many cases even below optical. This implies that synchrotron emission from these fast cooling electrons will be far above observational limits in the X-ray and in the optical. This problem is related, but not directly similar, to the low energy spectral index of the prompt $\\gamma$ emission which is much harder than the one predicted by fast cooling.    We consider three possible ways to overcome the excessive emission at low frequencies. The electrons can be accelerated before they cool down, they can escape rapidly from the large magnetic field zone or they may remain confined within sub-regions with low magnetic field where they are accelerated and then radiate less efficiently. The acceleration can be continuous \\cite{Ghisellini(1999),Kumar(2008)}. In this case the electrons are kept more or less at the same Lorentz factor during the whole emission episode posing a problem for production of the higher part of the energy spectrum. Alternatively the acceleration can be sporadic and the electrons are accelerated within ``re-acceleration sites\" (this is the situation in e.g. the ICMART model \\cite{Zhang(2011)}). The acceleration in these ``re-acceleration sites\" should be modest, by  a factor less than three in energy. This implies that this scenario requires two different acceleration process. First an initial acceleration that brings the electrons to $\\gamma_m \\sim 5000$ and then a second re-acceleration process that accelerates the electrons mildly. An additional constraint  is that all electrons present must be accelerated in the first acceleration phase, otherwise slow electrons would be accelerated to mild energies within the ``re-acceleration sites\" and will produce spurious low-energy synchrotron emission. Due to the re-acceleration only very few electrons are needed to produce the observed flux. Hence, only very few protons must be present in the outflow. Consequently all acceleration models require a very small baryonic loading (an extremely large $\\sigma$).   An alternative scenario is one in which the magnetic field is inhomogeneous and the electrons cool efficiently only in ``emission sites\" in which the magnetic field is stronger. However, this model requires that the magnetic field within the emission sites should be significantly larger than the field in the outer regions. This requires an external source of pressure that confines the magnetic field in the emission sites and it is incompatible with the situation which we consider in which the outflow is Poynting flux dominated. This might be applicable, however, in other situations in which the magnetic field is weaker.   Finally, motivated by reconnection models, electrons could accelerate and radiate in small sub-volumes of the emitting region with lower magnetic fields, where their cooling is less efficient. However, this model requires some extreme conditions. The magnetic fields within these regions should be very small (typical values of $\\epsilon_B \\sim 10^{-6}$) leading to larger Larmour radii compared with the skin depth. Moreover, given the small skin depth, the life-time of the low field regions is very likely to be much shorter than the dynamical time scale, and it is not clear that electrons can be confined sufficiently long in these low magnetic field regimes.   To summarize, we find that very specific conditions have to be satisfied if the emitting region is Poynting flux dominated. To be viable any model in which the emission is in such a region must demonstrate how these conditions are reached. To our knowledge there is no known physical model that indeed demonstrates how these conditions can be met. The simplest possible resolution of this conundrum is that the outflow is not Ponynting flux dominated in the emission region. This implies that Poynting flux dominated jets dissipate their magnetic energy before the emission zone. Interestingly very different considerations, based on examination of propagations of relativistic MHD jets within stellar envelope, have led  recently \\cite{Bromberg+14} to reach a similar conclusion, namely that Poynting flux dominated jets must dissipate deep in the interior of Collapsars in order to be compatible with observations.   We thank Jonathan Granot, Rodolfo Barniol Duran and Daniel Kagan for helpful discussions. This research was supported by the ERC advanced research grant ``GRBs'', by the  I-CORE (grant No 1829/12) and  by ISF-NSFC research grant."
    },
    "1402/1402.4922_arXiv.txt": {
        "abstract": "Scaling properties of nuclei- and photon-initiated air showers are examined in wide primary energy range ($10^{14}\\div 10^{22}$~eV) taking into account Landau-Pomeranchuk-Migdal and geomagnetic field effects. It is shown that the invariance in lateral distribution and also the universal dependence between shower age and root mean square radius of electron component exist up to the highest energies. The implications of shower universality for reanalysis, cross-checking and validation of the results of different experiments together with decreasing of the influence of hadronic model uncertainties are discussed in detail. ",
        "introduction": "The concept of air shower universality dating back to pioneering works~\\cite{RossiGreisen,NishimuraKamata,Greisen} is subjected to the remarkable rise of interest during the last decade~(see e.g.~\\cite{Capdevielle2005,Nerling:2005fj,Giller:2005qz,Gora:2006,Apel:2008hi,Lipari:2008td,Lafebre:2009en,Matthews,Yushkov:2010}). In a broad sense, the so-called universality is expressed in the similarity of the spectra of low-energy particles in air showers at the same development stage. Phenomenology, possible generalisations, deeper understanding of the origin and limitations of such properties are obviously important to unveil their potential in the achievement of the unambiguous physical interpretation of the results of comparisons between air shower observables and the appropriate theoretical calculations. Various parameterizations of the correlations between different air showers characteristics have been obtained on the basis of universality; attempts have been made to work out new experimental data processing techniques that could, for example, suppress or separate the influence of the hadronic interaction model and improve the reliability of the results.  In this paper we investigate perspectives of model-independent scaling approach for lateral distribution of electrons in nuclei- and photon-initiated air showers and also the universal dependence between shower age and root mean square radius of electron component for the improved analysis of the experimental data, particulary to extract information about the variations in elemental composition of cosmic rays in almost a model-independent way. ",
        "conclusions": "The universality in extensive air shower development expressed by means of scale-invariant lateral distribution of electrons and the relation between root mean square radius and depth of shower maximum is evidently retaining in a common form in extremely wide primary energy range (from $10^{14}$ to $10^{22}$~eV) neglecting possible exotic hadronic phenomena. This formalism could be useful in development of new effective techniques for reanalysis, cross-checking and validation of the results of different experiments together with decreasing of the influence of hadronic model uncertainties allowing to raise confidence in the results coming from the comparisons between air showers observables and theoretical calculations."
    },
    "1402/1402.6320_arXiv.txt": {
        "abstract": "We report the results from a 19-h integration with the SPIRE Fourier Transform Spectrometer aboard the \\emph{Herschel Space Observatory} which has revealed the presence of a molecular outflow from the Cosmic Eyelash (SMM\\,J2135\\(-\\)0102) via the detection of blueshifted OH absorption.  Detections of several fine-structure emission lines indicate low-excitation \\hii{} regions contribute strongly to the \\cii{} luminosity in this \\(z = 2.3\\) ULIRG.  The OH feature suggests a maximum wind velocity of \\(700 \\, \\kms\\), which is lower than the expected escape velocity of the host dark matter halo, \\({\\approx} 1000 \\, \\kms\\).  A large fraction of the available molecular gas could thus be converted into stars via a burst protracted by the resulting gas fountain, until an AGN-driven outflow can eject the remaining gas. ",
        "introduction": "\\label{introduction}  With only 10--30 per cent of the gas of a giant molecular cloud (GMC) undergoing conversion into stars, and another 40 per cent of that returning to the interstellar medium (ISM) through supernovae and stellar winds, star formation is an inefficient process \\citep{2011ApJ...729..133M}.  The rapid dispersal of GMCs by the combined radiation and mechanical pressure exerted by newly-formed massive stars is responsible for this inefficiency, lowering the intensity but extending the duration of the star formation resulting from an initial quantity of molecular material.  This is visible in the relatively small observed reduction in the gas available for star formation since \\(z \\sim 2\\) \\citep[of the order of \\(10 \\times\\); e.g.][]{2013ARA&A..51..105C}, and in the similar decrease in the cosmic star formation rate (SFR) in galaxies over the same period \\citep[e.g.][]{2013ApJ...770...57B, 2014MNRAS.437.3516S}.  Conversion to stars is not the only fate of molecular gas; stellar feedback, along with that from AGN, is also capable of expelling interstellar material from a galaxy entirely \\citep{2005ARA&A..43..769V}.  \\begin{figure*} \\centering \\includegraphics[]{figures/zoom.eps} \\caption{Left: rest-frame stack centred on the OH 119.233 \\um{} line at \\(z = 2.32591\\), the low-\\(J\\) luminosity-weighted redshift derived in \\citet{2011ApJ...742...11S}. The continuum flux level is normalized to that from the spectral energy distribution (SED) \\citep{2010A&A...518L..35I}. The vertical black dashed lines demarcate the region displayed in the right-hand plot.  Right: zoom-in on the OH absorption line with the rest-frame primary atomic and ionic fine-structure emission lines in the SPIRE FTS spectral range overplotted.  The lensing amplification as a function of velocity is shown as a dotted black line, suggesting a higher amplification for \\cii{} relative to the other emission lines.  Note that the SPIRE FTS produces a sinc line profile, with a width constant in frequency, and that all of the atomic and ionic lines are spectrally unresolved.  Vertical dashed lines are displayed at the velocities of the CO-bright clumps identified in \\citet{2013MNRAS.436.2793D}.  The vertical shaded region corresponds to the derived median and \\(1\\sigma\\) velocity of the OH feature assuming two identical components and corresponds to a shift of \\({-} 240 \\pm 110 \\,\\kms\\) relative to the CO clump `X', which contains the densest molecular gas, and which we associate with the primary source of the \\nii, \\oiii{} and \\oi{} emission (see \\S\\ref{sec:fs_lines}).} \\label{fig:zoom} \\end{figure*}  Occurring as galaxy-scale superwinds, these outflows are observed locally to be primarily powered by star formation, excluding those generated in late-stage galaxy mergers, which are typically driven by AGN \\citep{2013ApJ...776...27V}.  Superwinds seem to be more common and powerful at high redshift \\citep[e.g.][]{2009ApJ...692..187W, 2011MNRAS.418.1071B} where the majority of outflowing gas arises from strongly star-forming, high-mass galaxies and the outflow strength increases with SFR surface density \\(\\Sigma _\\text{SFR}\\) \\citep{2012ApJ...761...43N}.  Supernovae \\citep{2009ApJ...697.2030S}, stellar winds \\citep{2005ApJ...618..569M}, radiation pressure \\citep{2011ApJ...735...66M} and cosmic rays \\citep{2012MNRAS.423.2374U} may all contribute to the driving pressure.  While star formation is expected to produce wind velocities as high as 700 \\kms{} \\citep{2005ApJ...621..227M}, this may be insufficient for the material to escape the global potential well, leading to a `fountain': material subsequently falls back, leading to star formation over a time-scale much greater than that of the original burst \\citep{2013MNRAS.430.1901H}.  AGN may be required for complete expulsion of gas to the surrounding medium, since some seem able to drive outflow velocities of over 1000\\,\\kms{} \\citep{2005ApJ...632..751R, 2013ApJ...775..127S}.  As with the ISM from which it originates, a superwind is expected to comprise multiple gas phases moving with different velocities \\citep[e.g.][]{2005ApJ...632..751R, 2010A&A...518L.155F}.  Whilst observations of X-ray and UV absorption lines have indicated ionized and neutral atomic components \\citep[e.g.][]{2012ApJ...758..135K}, the presence of molecular gas may also be required (depending on the time-scale for material to cycle from atomic to molecular) if future star formation is to be reduced.  CO and OH are the most common non-H\\(_2\\) molecules at the temperatures and densities involved, with outflows inferred from observations of both species towards local galaxies and QSOs \\citep[e.g.][]{2009ApJ...700L.104S,2010A&A...518L..41F}, indicating a correlation between AGN luminosity and maximum outflow velocity \\citep{2013ApJ...775..127S, 2013ApJ...776...27V}, and showing that outflows are prevalent among local ULIRGs.  OH is particularly well suited for tracing molecular gas in galactic winds. It possesses a dipole moment \\({\\sim} 15 \\times\\) that of CO.  Due to the high densities required to thermalize OH rotational transitions \\citep[\\({\\sim} 10^9 \\, \\text{cm}^{-3}\\);][]{1981ApJ...244L..27S}, the strongest OH line is typically the ground-state absorption \\(^{16}\\text{OH} \\; ^2\\Pi_{3/2} \\; \\J{\\frac{3}{2}}{\\frac{5}{2}} \\), split by lambda-doubling into rest-frame 119.233 and 119.441\\,\\um{} transitions.  Back-lit molecular gas provides a prime opportunity to observe these transitions, with a blueshift expected from the radial movement of an outflow.  Despite the likely higher prevalence of superwinds at high redshift, and the potential advantages of observing OH, the far-infrared (FIR) rest-frame wavelengths of these transitions necessitate the use of space-based facilities. The consequential small apertures have meant that observations have been restricted to local galaxies \\citep[e.g.][]{2013ApJ...775..127S, 2013ApJ...776...27V}.  Previous studies of high-redshift outflows have generally utilized UV lines \\citep[e.g.][]{2010ApJ...717..289S,2012ApJ...760..127M}, which are generated within ionized material and thus insensitive to any molecular phase, critical to determination of whether superwinds or \\emph{in situ} star formation dominates the reduction in interstellar molecular gas and SFR density at early epochs.  Gravitational lensing provides a means to circumvent these limitations.  The galaxy, SMM\\,J2135\\(-\\)0102, hereafter \\smm, was identified in the field of a massive, lensing cluster via ground-based 870-\\um{} imaging \\citep{2010Natur.464..733S}.  Subsequent imaging with high spatial resolution constrained the lensing model, showing that \\smm{} is resolved into an \\({\\sim} 5 \\, \\text{arcsec}\\) long arc, mirrored about the critical curve, with a total magnification \\(\\mu = 37.5 \\pm 4.5\\) \\citep{2011ApJ...742...11S}. CO measurements of its velocity field reveal a disc with rotation speed \\(320 \\, \\kms\\), dominated by several star-forming clumps, each of radius 100--200\\,pc, distributed across a radius of \\({\\sim} 2 \\, \\text{kpc}\\).  Here, we describe the discovery of a molecular outflow from \\smm, as betrayed by the detection of the 119\\,\\um\\ OH transition in absorption during a number of long integrations with the \\emph{Herschel Space Observatory}.  Throughout, we use \\emph{Planck} 2013 cosmology \\citep{2013arXiv1303.5076P} with \\(H_0 = 67.8 \\, \\kms \\, \\text{Mpc}^{-1}\\), \\(\\Omega_\\text{m} = 0.307\\) and \\(\\Omega_\\Lambda = 0.693\\). ",
        "conclusions": "\\label{conclusions}  We present the first detection of OH absorption in a high-redshift galaxy, exploiting the SPIRE FTS aboard \\herschel, and interpret this as a starburst-driven molecular outflow.  We additionally detect emission lines of \\cii\\,158\\,\\um, \\nii\\,122\\,\\um, \\oiii\\,88\\,\\um{} and \\oi\\,63\\,\\um, finding line luminosities similar to local ULIRGs and indicative of a significant contribution to the \\cii{} luminosity from \\hii{} regions.  PDR models indicate \\(n\\) and \\(G_0\\) values of \\(10^{2 \\text{--} 3}\\) \\(\\text{cm}^{-3}\\) and \\(10^{2 \\text{--} 3}\\) Habing fields, respectively.  The 119-\\um{} OH absorption is blueshifted by \\(240 \\pm 110 \\, \\kms\\) relative to the star-forming clump, `X', pinpointed previously by low-\\(J\\) CO observations, and which we associate with the observed \\nii{} 122-\\um{} and \\oiii{} 88-\\um{} emission.  This molecular component to a starburst-driven wind, possesses a maximum velocity of \\(700 \\pm 200 \\, \\kms\\), smaller than the escape velocity expected for the dark matter halo.  A large fraction of the available molecular gas could thus be converted into stars via a burst protracted by the resulting gas fountain, until an AGN-driven outflow can eject the remaining gas."
    },
    "1402/1402.5924_arXiv.txt": {
        "abstract": "As the second part of our study, in this paper we proceed to explore the dynamics of the high-inclination Twotinos in the 1:2 Neptune mean motion resonance (NMMR). Depending on the inclination $i$, we show the existence of two critical eccentricities $e_a(i)$ and $e_c(i)$, which are lower limits of the eccentricity $e$ for the resonant angle $\\sigma$ to exhibit libration and asymmetric libration, respectively. Accordingly, we have determined the libration centres $\\sigma_0$ for inclined orbits, which are strongly dependent on $i$. With initial $\\sigma=\\sigma_0$ on a fine grid of $(e, i)$, the stability of orbits in the 1:2 NMMR is probed by 4-Gyr integrations. It is shown that symmetric librators are totally unstable for $i\\ge30^{\\circ}$; while stable asymmetric librators exist for $i$ up to $90^{\\circ}$.  We further investigate the 1:2 NMMR capture and retention of planetesimals with initial inclinations $i_0\\le90^{\\circ}$ in the planet migration model using a time-scale of $2\\times10^7$ yr. We find that: (1) the capture efficiency of the 1:2 NMMR decreases drastically with the increase of $i_0$, and it goes to 0 when $i_0\\gtrsim60^{\\circ}$; (2) the probability of discovering Twotinos with $i>25^{\\circ}$, beyond observed values, is roughly estimated to be $\\lesssim0.1$ per cent; (3) more particles are captured into the leading rather than the trailing asymmetric resonance for $i_0\\le10^{\\circ}$, but this number difference appears to be the opposite at $i_0=20^{\\circ}$ and is continuously varying for even larger $i_0$; (4) captured Twotinos residing in the trailing resonance or having $i>15^{\\circ}$ are practically outside the Kozai mechanism, like currently observed samples. ",
        "introduction": "The Kuiper belt objects (KBOs) are icy celestial bodies near and beyond the orbit of Neptune in the outer Solar system. One of the most outstanding mysteries of these distant objects is the unexpected high orbital inclination ($i$). In the previous work (Li et al. 2014, hereafter LZS14), we began to study this issue by considering the dynamics of the high-inclination Plutinos. By definition, a Plutino occupies the 2:3 Neptune mean motion resonance (NMMR) at semimajor axis $a\\approx39.4$ au. We have defined the special libration centre (SLC) of the resonant angle $\\sigma_{2:3}$ at the minimum of the averaged disturbing function $R({\\sigma_{2:3}})$, which corresponds to a stable equilibrium point because the variation of the semimajor axis $da/dt(\\propto{\\partial R}/{\\partial \\sigma_{2:3}})=0$. Then we designated the mean value of $\\sigma_{2:3}$ during the time evolution as the general libration centre (GLC) at $180^{\\circ}$, i.e., the usually called ``libration centre''. For low-inclination orbits, the position of the SLC is fixed exactly at the GLC. But for high-inclination orbits, the SLC becomes strongly dependent on the argument of perihelion $\\omega$ and oscillates around the GLC with large amplitude (Gallardo 2006). Nevertheless, the average value of the SLC over $0^{\\circ}\\le \\omega \\le360^{\\circ}$ always matches the GLC. This outcome suggests a new method that can help us to easily determine the libration centre of the other NMMRs for any inclined orbits. We refer the reader to LZS14 for more details on the calculations of the SLC.  In LZS14 we also updated the possible $i$-range of the potential Plutinos.  Using the $N$-body simulations, we first investigated the stability of high-inclination candidates populating in the 2:3 NMMR for times up to the age of the Solar system. The dynamical map has been built up on the initial $(a, i)$ plane at different eccentricities ($e$). We have found that the stable resonant orbits could cover the whole inclination space of $i\\le90^{\\circ}$. Based on this result, we further explored the outward transportation of planetesimals with initial $i$ up to $90^{\\circ}$ in the framework of the planet migration and 2:3 NMMR sweeping model (Malhotra 1995). Our results showed that the resonant capture and retainment is allowed for any inclined or even perpendicular orbits. Beside the aforementioned main concerns, the role of high $i$ in both the formation of Plutinos during Neptune's migration and the later long-term evolution has been discussed in detail.  \\begin{figure} \\hspace{-0.7cm} \\centering \\includegraphics[width=9cm]{Fig1.JPG} \\caption{Distribution of eccentricities and inclinations for the currently observed Twotinos (as of October 2012). The symmetric librators are stars, the leading librators are upward triangles, and the trailing librators are downward triangles. The filled symbols refer to Twotinos experiencing the Kozai mechanism (see Section 3.2.2 for definition).} \\label{observed} \\end{figure}  As the continuation of our study of the high-inclination KBOs, in this paper we investigate a population of bodies with $a\\approx47.8$ au in the 1:2 NMMR, which are known as Twotinos. It is of further interest to present a global view of the dynamics of Twotinos on inclined orbits. In the MPC (Minor Planet Center) database\\footnote{http://www.minorplanetcenter.net/iau/lists/TNOs.html}, about 24 Twotinos have been registered and their osculating $e$ and $i$ at epoch 2012 September 30 are shown in Fig. \\ref{observed}. For the sake of identifying these Twotinos, we numerically integrated the trajectories of observed KBOs with $a=47.8\\pm0.8$ au for a timespan of $10^8$ yr, under the gravitational perturbations of four Jovian planets. Then we examined the time evolution of the resonant angle $\\sigma=\\lambda_N-2\\lambda+\\varpi$ corresponding to the 1:2 NMMR, where $\\lambda$ and $\\varpi$ are the mean longitude and the longitude of perihelion of the KBO, respectively, and $\\lambda_N$ is the mean longitude of Neptune. In this way a KBO is regarded as a Twotino if it exhibits libration of $\\sigma$, i.e., the resonant amplitude $A_{\\sigma}<180^{\\circ}$. It is seen that the Twotinos possess high inclinations up to about $25^{\\circ}$, which seems to be in accord with the stable $i$-range obtained in previous literatures (Melita \\& Brunini 2000; Nesvorn\\'{y} \\& Roig 2001; Lykawka \\& Mukai 2007; Tiscareno \\& Malhotra 2009; Gladman et al. 2012). In spite of this, we are still eager to know whether $i\\sim25^{\\circ}$ is actually the limit of all the potential Twotinos. It must be noted that here we just confirmed or rejected the 1:2 NMMR librators with nominal orbital elements given in the MPC site. If one takes into account the uncertainties of the orbital elements, there might yield some other librators near 47.8 au, and particularly a part of them could have $i>25^{\\circ}$. Besides, observational selection effects would also work against the detection of high-inclination objects.  It is well known that, for the type of 1:$n$ exterior mean motion resonances, there are three possible resonant modes: the usual {\\it symmetric} libration of $\\sigma$ around $180^{\\circ}$; and two separate {\\it asymmetric} librations for $0^{\\circ}<\\sigma<180^{\\circ}$ (leading) and for $180^{\\circ}<\\sigma<360^{\\circ}$ (trailing) (Beaug\\'{e} 1994; Morbidelli et al. 1995).  The leading and trailing librators reach their perihelia at longitudes ahead of and behind Neptune's longitude, respectively. In the framework of the planar circular restricted 3-body problem, Malhotra (1996) pointed out that the libration centre ${\\sigma}_0$ is not a well-defined one but depends upon the orbital eccentricity $e$ of the particle. Later, the value of ${\\sigma}_0$ related to the asymmetric 1:2 NMMR as a function of $e$ was estimated by Nesvorn\\'{y} \\& Roig (2001). However, the dependence of ${\\sigma}_0$ on the inclinations of Twotinos has not been researched until now. This is due to the fact that the analytical or semi-analytical computations of asymmetric periodic orbits are not so easy in the three-dimensional restricted 3-body problem (Kotoulas 2005).  With respect to the history of Neptune's outward migration, the relative populations of the leading and trailing Twotinos may be an important diagnostic for the migration time-scale $\\tau$. Chiang \\& Jordan (2002) and Murray-Clay \\& Chiang (2005) proposed that planetesimals are less likely to be trapped into the leading rather than the trailing 1:2 NMMR for $\\tau$ shorter than $\\sim10^7$ yr. But according to observational data shown in Fig. \\ref{observed}, 10 Twotinos are in the leading resonance and only 4 are in the trailing resonance, such a number difference may support an even larger $\\tau$ (Lykawka \\& Mukai 2007). Indeed, a time-scale of $\\tau=2\\times10^7$ yr was derived by Li et al. (2011) from $N$-body simulations for the orbital evolution of Jovian planets embedded in a self-gravitating planetesimal disc. This reasonable value has already been adopted in LZS14.  Moreover, Murray-Clay \\& Chiang (2005) also discussed the capture into the asymmetric 1:2 NMMR influenced by planetesimal's initial eccentricity $e_0$ and initial semimajor axis $a_0$. Given the time-scale of $\\tau=10^7$ yr and $a_0\\ge37.7$ au (1 au exterior to the initial 1:2 NMMR), the ratio of leading to trailing particles can be infinitely small for $e_0=0.01$; while it grows with lager $e_0$ and would nearly approach 1 for $e_0=0.05$, i.e., the difference in populations vanishes. As we argued in LZS14, the primordial planetesimals may have high initial inclinations $i_0$ prior to the onset of the resonance sweeping. Therefore, in order to complete a survey of the 1:2 NMMR capture in the entire orbital element space, a detailed analysis of the role of $i_0$ is clearly warranted.  The rest of this paper is organized as follows. In Section 2, we reveal the shift of the libration centre ${\\sigma}_0$ due to high $i$, and the complex resonant behavior in the 1:2 NMMR. In Section 3, we probe the long-term stability of high-inclination objects in the 1:2 NMMR, and show that $i$ of stable librators could be as high as $90^{\\circ}$. Here, the $i$-dependent ${\\sigma}_0$ is crucial to choose the initial conditions. In Section 4, for different initial inclinations $i_0$, we discuss in detail the resonance capture and orbital evolution of planetesimals in the planet migration model using a time-scale of $\\tau=2\\times10^7$ yr. In Section 5, we conclude this paper with a summary of our main findings, and discuss the theoretical models and current observations.   \\section[]{The resonant behavior and the libration centre}  \\begin{figure} \\hspace{-0.7cm} \\centering \\includegraphics[width=9cm]{Fig2.JPG} \\caption{The phase space of the 1:2 NMMR for eccentricity $e=0.15>e_c$ in the framework of the planar circular restricted 3-body problem. The horizontal dashed line at 47.8 au denotes the location of nominal 1:2 NMMR. The three vertical dashed lines indicate the symmetric libration centre at $180^{\\circ}$, and the asymmetric ones at $90^{\\circ}$ (leading) and $270^{\\circ}$ (trailing).} \\label{phase} \\end{figure}  Restricted to the planar case ($i=0^{\\circ}$), the dependence of the libration centres of the 1:2 NMMR on the eccentricity $e$ has been published for many years and is known perfectly well (Beaug\\'{e} 1994; Morbidelli et al. 1995; Malhotra 1996). At small values of $e$, the phase space of the 1:2 NMMR is largely regular like a pendulum system, and there is only a symmetric libration centre $\\sigma^{\\pi}_0$ at $180^{\\circ}$. While for $e$ exceeding a critical value $e_c\\sim 0.04$, the typical $(\\sigma, a)$ phase space of this resonance is sketched in Fig. \\ref{phase}. The symmetric resonance island splits into two asymmetric resonance islands, where the resonant amplitudes $A_{\\sigma}$ have an upper bound ($<90^{\\circ}$). Meanwhile, the symmetrically librating orbits with larger $A_{\\sigma}$ (``horseshoe'' orbits) can still exist, surrounding the asymmetrically librating orbits (``tadpole'' orbits). Note that the locations of asymmetric centres can be modified by the value of $e$. For $e=0.15>e_c$, as shown in Fig. \\ref{phase}, the leading centre $\\sigma^L_0$ and the trailing centre $\\sigma^T_0$ are placed at $90^{\\circ}$ and $270^{\\circ}$, respectively.   \\begin{figure} \\centering \\begin{minipage}[c]{0.6\\textwidth} \\centering \\hspace{-3cm} \\includegraphics[width=9cm]{Fig3a.JPG} \\end{minipage} \\begin{minipage}[c]{0.6\\textwidth} \\centering \\hspace{-3cm} \\includegraphics[width=9cm]{Fig3b.JPG} \\end{minipage} \\caption{The behavior of the SLC for the 1:2 NMMR for (a) $i=10^{\\circ}$ and (b) $i=30^{\\circ}$. The upper, middle and lower panels correspond to the cases of $e<e_a$, $e_a<e<e_c$ and $e>e_c$, respectively. In the middle panel of graph (b), notice at $\\omega=270^{\\circ}$ that the SLC is approximately $34 ^{\\circ} $ away from $\\pm180^{\\circ}$.} \\label{SLC} \\end{figure}   To evaluate the libration centres of the 1:2 NMMR for inclined orbits, we followed the semi-analytical method developed in LZS14. The libration centre is denoted by the GLC, i.e., ${\\sigma}_0=({\\sigma_{max}+{\\sigma}_{min}})/2$, which can be simply and precisely calculated as the mean value of the SLC. And for the resonant amplitude $A_{\\sigma}=\\sigma_{max}-{\\sigma}_0$, it has a lower limit that is virtually the amplitude of the SLC. Unlike in the planar case, for a specific $e$, the SLC would be no longer fixed at the GLC but moves right along with the argument of perihelion $\\omega$ when the non-zero $i$ has been introduced (Gallardo 2006).  The general characteristics of the 1:2 NMMR can be understood by analyzing the behavior of the SLC for a series of $\\omega$ between $0^{\\circ}$ and $360^{\\circ}$, as presented in Fig. \\ref{SLC}:  \\begin{enumerate} \\item If $e<e_a$, the motion of the SLC is unbounded (Fig. \\ref{SLC}, upper panels). Since the amplitude of the SLC is the lower limit of $A_{\\sigma}$, this implies that the resonant angle $\\sigma$ would eventually take all values between $0^{\\circ}$ and $360^{\\circ}$.  As we will see in the next section, the resonant angle $\\sigma$ would be simply circulating in this case; while a plausible scenario where there is a libration of $\\sigma$ around SLC but SLC itself circulates would never occur. Hence we consider the critical eccentricity $e_a$ to correspond to the lower limit of $e$ for librations in the 1:2 NMMR.  \\item If $e_a<e<e_c$, then the motion of the SLC becomes bounded. For relatively small $i$ (e.g., Fig. \\ref{SLC}a, middle panel), the SLC does not pass by the location of $0^{\\circ}$, and shows libration around the GLC at $180^{\\circ}$. But for relatively large $i$,  it is found that the SLC never achieves $180^{\\circ}$ and then the GLC changes to $0^{\\circ}$. This is illustrated by the middle panel of Fig. \\ref{SLC}b, where the interval of the vertical axis is adopted to be [-$180^{\\circ}$, $180^{\\circ}$] for a simpler visualization. Based on our calculations, the shift of the GLC from $180^{\\circ}$ to $0^{\\circ}$ takes place at a tentative inclination of $\\sim15^{\\circ}$. Note that the resonant orbits with small $A_{\\sigma}$ ($<90^{\\circ}$) are permitted here.\\\\ For the perihelic conjunctions at $\\sigma=0^{\\circ}$, the vertical distance from a small-$A_{\\sigma}$ Twotino to Neptune's orbital plane is roughly estimated by $d=a(1-e_c)\\sin{i}$. This yields a value of $d\\approx11.4$ au for $i=15^{\\circ}$, and this distance is monotonically increasing with $i$. Therefore, in the absence of strong gravitational perturbation from Neptune, the libration of $\\sigma$ around $0^{\\circ}$ could be a possibly stable configuration.  \\item If $e>e_c$, the motion of the SLC is restricted to either the interval of $[\\sigma_1, \\sigma_2]\\subset (0^{\\circ}, 180^{\\circ})$  or $[\\sigma_3, \\sigma_4] \\subset (180^{\\circ}, 360^{\\circ})$ (Fig. \\ref{SLC}, lower panels). This implies the appearance of asymmetric islands in the 1:2 NMMR (see Fig. \\ref{phase}), and there are two GLCs: $\\sigma^L_0=(\\sigma_1+\\sigma_2)/2$ and $\\sigma^T_0=(\\sigma_3+\\sigma_4)/2$. As mentioned earlier, $e_c$ represents the low limit for $e$ for the permission of the asymmetric 1:2 NMMR. It is necessary to stress that, our SLCs are determined at the minimum of the resonant disturbing function, but they do not exclude the existence of large-$A_{\\sigma}$ symmetric librations around the GLC $\\sigma^{\\pi}_0=180^{\\circ}$ (Gallardo 2006). \\end{enumerate}  \\begin{figure} \\hspace{-0.7cm} \\centering \\includegraphics[width=9cm]{Fig4.JPG} \\caption{The critical eccentricities $e_a$ (dashed line) and $e_c$ (solid line) for the 1:2 NMMR as functions of the inclination $i$. Note that these two quantities both increase with $i$.} \\label{eaec} \\end{figure}  The two critical eccentricities $e_a$ and $e_c$ for the 1:2 NMMR are both functions of the inclination $i$. Fig. \\ref{eaec} gives the values of $e_a(i)$ and $e_c(i)$ derived from the different behaviors of the SLCs for orbits with $i$ from $0^{\\circ}$ to $90^{\\circ}$. For the planar case of $i=0^{\\circ}$, we have $e_a=0$ and $e_c=0.037$, which are in reasonably good agreement with those obtained by Beaug\\'{e} (1994) and Malhotra (1996). With increasing $i$, these two quantities both shift to larger values, and $e_c$ reaches its peak of $\\sim0.21$ when $i$ exceeds about $60^{\\circ}$. Especially, for the highest inclination of $\\sim25^{\\circ}$ observed in Fig. \\ref{observed}, the corresponding $e_c$ is about 0.12.  This may agree with the fact that real Twotinos preferentially have large eccentricities. The only exception is a Twotino with $e\\sim0.07$ in the bottom-left corner of Fig. \\ref{observed}. Due to the nearly zero inclination, it has sufficiently large $e$ relative to $e_c(i=0^{\\circ})=0.037$.  \\begin{figure} \\hspace{-0.7cm} \\centering \\includegraphics[width=9cm]{Fig5.JPG} \\caption{The locations of the GLC with respect to inclination for the 1:2 NMMR, at five eccentricities of $e=0.05$ (black), 0.15 (red), 0.25 (blue), 0.35 (yellow) and 0.40 (green).} \\label{GLC} \\end{figure}  Taking advantage of the above calculations, we present in Fig. \\ref{GLC} the locations of the GLC with respect to inclination for the 1:2 NMMR, at five representative eccentricities of $e=0.05$, 0.15, 0.25, 0.35 and 0.40 (covering the observational range). The resonant states of Twotinos can be divided into the following three typical categories:  \\begin{enumerate} \\item For $e=0.05$, the asymmetric libration has ceased to exist when $i>5^{\\circ}$, while only the symmetric libration around $180^{\\circ}$ is possible. Furthermore, if $i$ continues to grow beyond about $12^{\\circ}$, objects would be surely outside the 1:2 NMMR since the libration of $\\sigma$ is prohibited by the circulation of the SLC.  \\item For $e=0.15$, the asymmetric libration is allowed when $i\\le35^{\\circ}$. We find that the libration centre $\\sigma^L_0$ ($\\sigma^T_0$) is at $90^{\\circ}$ ($270^{\\circ}$) for $i=0^{\\circ}$, exactly as that shown in Fig. \\ref{phase}, and it decreases (increases) continuously to $73^{\\circ}$ ($287^{\\circ}$) as $i$ increases to $35^{\\circ}$. But for even larger $i$, equivalent to the case of $e_a<e<e_c$, it appears an uncommon libration centre at  $0^{\\circ}$, which will be further discussed in a later section.  \\item For $e\\ge0.25$, the asymmetric libration is possible for any value of $i$, because $e$ is larger than the maximum $e_c(i)\\sim0.21$. Along with the increase of $i$, the magnitude of variation in $\\sigma^L_0$ and $\\sigma^T_0$ is either above or near $10^{\\circ}$.  \\end{enumerate}  For the numerical investigation of the long-term stability of high-inclination Twotinos, our semi-analytical predictions could provide a wealth of information on the choice of initial conditions. Accordingly, for a given $e$, we can set the initial resonant angle $\\sigma$ of all orbits to be the exact GLC value within a proper $i$-range. ",
        "conclusions": "To continue our previous study of the high-inclination KBOs, in which we explored the dynamic of Plutinos with $i$ up to $90^{\\circ}$ (in LZS14), this second paper is devoted to Twotinos. We are acutely aware of the existence of symmetric and asymmetric motions in the 1:2 NMMR, thus Twotinos may furnish a unique window on the distant past of the outer Solar system.  Using the semi-analytical method developed in LZS14, for the first time, we provide a global picture of the dynamical aspects of the 1:2 NMMR in three dimensions. Depending on the critical eccentricities $e_a(i)$ and $e_c(i)$ as functions of the inclination $i$, three different behaviors of the resonant angle $\\sigma$ have been found: (1)  if $e<e_a$, the circulation of $\\sigma$ by $360^{\\circ}$ is certain; (2) if $e_a<e<e_c$, there appear a libration of $\\sigma$ around $180^{\\circ}$ for $i<15^{\\circ}$, and another libration around $0^{\\circ}$ for larger $i$; and (3) if $e>e_c$, the asymmetric libration of $\\sigma$ becomes possible, meanwhile the symmetric libration around $180^{\\circ}$ is also allowed. At zero inclination, we have $e_a=0$ and $e_c=0.037$, equally consistent with earlier works (Beaug\\'{e} 1994; Malhotra 1996). But with increasing $i$, these two quantities both shift to larger values, and $e_c$ achieves a maximum of $\\sim0.21$ when $i>60^{\\circ}$. These characteristics indicate that, for highly inclined orbits, the asymmetric libration islands can only emerge at large eccentricities. Accordingly, we calculated the libration centres $\\sigma_0$ for 1:2 resonant orbits with $0^{\\circ}\\le i \\le 90^{\\circ}$, at specific $e$. We have shown that, for high-inclination Twotinos, their $\\sigma_0$ are strongly dependent on $i$ and could be very different from those of low-inclination ones.  With the initial $\\sigma=\\sigma_0$ on a particular grid of $(e, i)$ based on the above semi-analytical predictions, test particles initialized on 1:2 resonant orbits were numerically integrated with four Jovian planets for times up to the age of the Solar system. Our principal findings for the long-term stability of Twotinos can be briefly summarized as follows.  (1) In the symmetric resonance, we find that the stable horseshoe orbits have $e\\le0.25$, $i\\le20^{\\circ}$ and resonant amplitudes $A_{\\sigma}<175^{\\circ}$. The most regular motion only exists in the low inclination region of $i\\lesssim5^{\\circ}$, where initial horseshoe orbits can exhibit persistent libration of $\\sigma$ throughout the entire 4 Gyr simulation. For more inclined particles, they are fated to undergo the frequent alternation between libration and circulation during the evolution, while many of them can survive as symmetric librators in the 1:2 NMMR.  (2) In the asymmetric resonance, we find that the inclination coverage of stable tadpole orbits is $\\le30^{\\circ}$ for moderate $e=0.15$, and as high as $90^{\\circ}$ for large $e$ between 0.25 and 0.4. By measuring the full resonant amplitude $A^f_{\\sigma}$ over the 4 Gyr interval, the dynamical maps have been constructed on the plane of initial resonant amplitude versus initial inclination. It is shown that the stable regions are generally delimited by initial $A_{\\sigma}<50^{\\circ}$, but their widths are quite distinct at different initial $i$. The unstable tadpole librators are principally accounted for by the appearance of the Kozai mechanism. However, their escaping routes from the 1:2 NMMR could be drastically different for small ($\\sim0^{\\circ}$), medium ($10^{\\circ}$--$20^{\\circ}$) and high ($30^{\\circ}$--$90^{\\circ}$) inclinations. Overall, there is no statistically significant discrepancy in the dynamical structure between the leading and trailing asymmetric resonances.  Then, we further investigated the evolution of primordially inclined planetesimals with respect to the 1:2 NMMR sweeping and capture, for which Neptune starts at $\\sim23.2$ au and migrates slowly with a time-scale of $\\tau=2\\times10^7$ yr. At each $i_0$ between $0^{\\circ}$ and $90^{\\circ}$ ($\\Delta i_0=10^{\\circ}$), 911 test particles with $37.7\\le a \\le 46.8$ au ($\\Delta a=0.01$ au) and $e=0.01$ have been introduced. Of surviving objects in the 1:2 NMMR zone at the end of the migration, according to the above stability analysis, we defined captured Twotinos as:  captured symmetric Twotinos (CSTs) with $A_{\\sigma}<175^{\\circ}$; captured leading Twotinos (CLTs) and captured trailing Twotinos (CTTs) with $A_{\\sigma}<50^{\\circ}$.  Our results show that the capture efficiency of the symmetric 1:2 NMMR, $R_S$, decays from 4.7\\% to 0.2\\% as the initial inclination $i_0$ increases from $0^{\\circ}$ to $30^{\\circ}$. While for $i_0\\ge40^{\\circ}$, there is a complete absence of CSTs. The final inclinations of CSTs suggest an upper limit of $i\\sim28^{\\circ}$ for potential symmetric Twotinos, which is just inside the border of the totally unstable horseshoe region at $i\\ge30^{\\circ}$.  As for the asymmetric 1:2 NMMR, the capture efficiency $R_L+R_T$ fails from 49.4 per cent for $i_0=0^{\\circ}$ to 0.4 per cent for $i_0=50^{\\circ}$. Particles originated in this $i_0$-range can possess small enough $A_{\\sigma}$ to reside deeply in the asymmetric islands, and therefore form long-lived residents. For few CLTs and CTTs with relatively large $A_{\\sigma}$ from the cases of $i_0\\ge60^{\\circ}$, they are actually embedded in or very close to unstable regions and will ultimately escape from this resonance. By combining the inclination distributions of CSTs, CLTs and CTTs, resonance sweeping capture may exclude the existence of Twotinos with $i>50^{\\circ}$. Moreover, this very mechanism predicts that the probability of discovering Twotinos with higher inclinations $i>25^{\\circ}$ than observed values, near the ecliptic plane, is only $\\sim0.1$ per cent.  Analytic arguments tell us that the ratio of CLTs to CTTs is largely determined by two factors: Neptune's migration time-scale and initial particle eccentricity (Murray-Clay \\& Chiang 2005). In addition, we found that the influence of initial particle inclination $i_0$ could also be great on this ratio, provided small particle eccentricity of $e\\sim0.01$. When $i_0\\le10^{\\circ}$, a particle is captured first into the symmetric 1:2 NMMR when Neptune migrates outward, then it will follow the dynamical pathway described in Murray-Clay \\& Chiang (2005). Due to the long migration time-scale of $\\tau=2\\times10^7$ yr, such symmetric librator is preferentially caught into the leading rather than the trailing asymmetric resonance. However, for the case of $i_0\\ge20^{\\circ}$, a particle captured by the sweeping 1:2 NNMR from a nearly circular orbit would never firstly fall into the symmetric resonance. This is because the libration of $\\sigma$ around $180^{\\circ}$ is not permitted for orbits with $i>15^{\\circ}$ and $e<e_c(i)$, as we have learned in the semi-analysis part. As a result, in the subsequent evolution, the probability of the high-inclination captive to enter the leading or trailing islands is not certain but would be continuously varying with $20^{\\circ} \\le i_0 \\le 50^{\\circ}$.  Among captured Twotinos from our simulations, we obtained a total of 4 potential $\\omega$-librators: 1 in the symmetric resonance, 3 in the leading asymmetric resonance. These Kozai Twotinos all come from the cases of $i_0\\le10^{\\circ}$. Despite somewhat excitation due to the Kozai mechanism, their $i$ would always undergo low value episode below $15^{\\circ}$ and end at $<12^{\\circ}$. This outcome seems compatible with today's observed void of Kozai Twotinos residing in the trailing islands or having $i>15^{\\circ}$.  Nevertheless, the planet migration model suffers from two major problems: (1) the fraction of captured Twotinos with high inclinations is 1--2 orders of magnitude too low to account for present observations; (2) the capture efficiency of the Kozai mechanism that we obtain is only $\\sim0.4$ per cent, whereas the observed value is $\\sim8.3$ per cent. These discrepancies argue that an unknown fraction of currently observed Twotinos is not likely to originate by the resonance sweeping mechanism from a planetesimal disk interior to 47.8 au. These particular objects indicate that there is a distinct source region for the high-inclination Twotinos, and they might be able to promise new insight into the dynamical history of the Kuiper belt. We expect our future works could tell."
    },
    "1402/1402.0506_arXiv.txt": {
        "abstract": "I discuss fluctuations in the neutral hydrogen density of the $z \\approx 2.3$ intergalactic medium and show that their relation to cosmic overdensity is strongly scale-dependent. This behaviour arises from a linearized version of the well-known ``proximity effect'', in which bright sources suppress atomic hydrogen density. Using a novel, systematic and detailed linear-theory radiative transfer calculation, I demonstrate how \\hi density consequently anti-correlates with total matter density when averaged on scales exceeding the Lyman-limit mean-free-path.  The radiative transfer thumbprint is highly distinctive and should be measurable in the Lyman-$\\alpha$ forest. Effects extend to sufficiently small scales to generate significant distortion of the correlation function shape around the baryon acoustic oscillation peak, although the peak location shifts only by $1.2$ percent for a mean source bias of $b_j=3$. The distortion changes significantly with $b_j$ and other astrophysical parameters; measuring it should provide a helpful observational constraint on the nature of ionizing photon sources in the near future. ",
        "introduction": "The Lyman-$\\alpha$ forest \\cite{Rauch98ARAA} is the imprint of the intergalactic medium (IGM) -- specifically, neutral hydrogen -- on the spectra of distant quasars.  At high redshift the rapidly-changing forest probes hydrogen reionization \\cite{Becker01Reionization,Fan02Reionization,2006ARA&A..44..415F}; at lower redshift, the forest has a steadier ionization state and is used to trace overall matter density fluctuations \\cite{Croft02,McDonald03,BOSS_Slosar11}. Correlating Lyman-$\\alpha$ fluctuations over small scales therefore places a strong constraint on modifications to the standard cold dark matter picture of structure formation \\cite{Viel_WDM_Lya_2005,Viel_WDM_Lya_2008,Boyarsky_WDM_Lya_2009}. More recently attention has turned to the large-scale forest's ability to constrain the baryon acoustic oscillation peak, providing an independent distance measurement for constraining dark energy \\cite{Slosar09BAO,BOSS_Slosar13,FontRibera_BOSS_QSO_Lya_X_BAO_2013}. In addition to probing the power spectrum in these ways, the observed $z<5$ forest constrains the thermal state of the intergalactic medium \\cite{Becker07,Bolton08}, allowing various interesting processes to be studied (such as helium reionization \\cite{Becker11HeIIreionization}).  When considering the forest after reionization, it is standard practice \\cite{Croft98,McDonald00,McDonald03} to model the IGM ionization state in the presence of a uniform background of ultraviolet (UV) photons.  Direct constraints on the Lyman-$\\alpha$ cloud temperatures \\cite{Bolton13IGMtemp} dictate that collisional ionization is unimportant except in systems that are dense enough to be substantially self-shielded from the radiation.  However the assumption that the UV background is uniform is known to be incorrect, since the constituent photons are actually generated by galaxies and quasars. One can distinguish two limits in which the approximation fails. First, on small scales, quasars are rare; depending on the fraction of photons they contribute (likely around $50\\%$ for $2<z<3$ \\cite{Faucher-Giguere09_UVB,HM12}) they can add significant shot noise on small scales.  Further fluctuations are imprinted by intrinsic variability in the IGM opacity \\cite{Maselli05_LyA_transmissivity}.  This and related astrophysical effects have been widely investigated elsewhere \\cite{Zuo92,Kollmeier03_LyA_LBG_connection,Meiskin04_LyA_radiation,Kollmeier06_LyA_galwinds,Viel_WDM_Lya_2008,Slosar09BAO,Mesinger09_LyA_fluctuating_reionisation,White10_LyA_sims,FontRibera_BOSS_QSO_Lya_X_2013} with the conclusion that, if properly accounted for, the added noise is not problematic for observational cosmology at $z<5$. Measurements at higher redshift, during the epoch of reionization, will be affected more strongly \\cite{Meiskin04_LyA_radiation,Mesinger09_LyA_fluctuating_reionisation} as the UV undulation amplitude increases.  In this paper I will consider the post-reionization IGM and place more emphasis on a second failure of the uniform-radiation assumption. This appears only when source clustering is taken into account on scales around the mean-free-path of an ionizing photon. By definition, regions separated by greater distances cannot efficiently exchange UV radiation. Ionization equilibrium will therefore depend on the density of sources in the local region; the bias of the forest on the largest scales will depend on the clustering of UV sources \\cite{Croft04_LyA_large_radiation,McDonald05,McQuinn11_LyA_fluctuations,White10_LyA_sims}.  This effect has received less attention to date, probably because the relevant scale is seemingly extremely large (the mean-free-path is of order $500\\,\\Mpc$ in comoving units \\cite{Rudie13MFP} at $z=2.4$). In fact, once redshifting and volume dilution are accounted for, the transition scale is somewhat smaller (more like $350\\,\\Mpc$ comoving; see Section \\ref{sec:absorption-terms}). To fully model such scales would require exceptionally large radiative-transfer simulations, with box sizes exceeding a gigaparsec to properly probe long-wavelength fluctuations.  \\updated{To achieve this, previous work has employed a combination of large dark-matter-only boxes and smaller hydrodynamic simulations \\cite{Croft04_LyA_large_radiation} or semi-analytic prescriptions \\cite{McDonald05}. In the former case the author reported a significant drop in large-scale flux power out to scales of $k^{-1} \\sim 70\\, h^{-1} \\,\\Mpc$ relative to the homogeneous-radiation control case. Even so, the result has not received widespread attention. This is likely because extending such state-of-the-art work consumes a great deal of computer time and, furthermore, appropriate empirical constraints for such large separations have seemed out of reach.}  The observational situation has now been changed radically by the BOSS (Baryonic Oscillation Spectroscopic Survey) project \\cite{2013Dawson_BOSS_overview}. The team have released results demonstrating the viability of measuring the correlation function of Lyman alpha clouds on large scales \\cite{BOSS_Slosar11,2013_Busca_BOSS_Lya_BAO,BOSS_Slosar13}. The major goal of BOSS is to measure the baryonic acoustic oscillation (BAO) feature in the correlation function at $100\\,h^{-1}\\,\\Mpc$ comoving. This is not so far off the reduced mean-free-path scale discussed above and derived in Section~\\ref{sec:absorption-terms}. It is timely, therefore, to reconsider the impact of large-scale fluctuations in the UV source density on the Lyman-$\\alpha$ forest.  The remainder of this paper proceeds systematically from first principles to a detailed linear-theory calculation of these effects. This should be highly complementary to numerical studies, and motivate further work in the area. I will ignore observational questions such as the transformation from \\hi to flux power spectrum, redshift-space distortions, redshift evolution and flux calibration biases -- since these require major computational machinery in themselves \\cite{2013_Busca_BOSS_Lya_BAO,BOSS_Slosar13} -- and focus on the bias of the physical intergalactic \\hi density at a single, fixed redshift. The quantitative results will be presented for $z=2.3$, around the mean redshift of observed Lyman-$\\alpha$ clouds \\cite{2013_Busca_BOSS_Lya_BAO}.  The approximations that allow this calculation to be completed are {\\it (i)} that the spatial variations in the UV spectrum are less important for \\hi than the changes in intensity (a `monochromatic approximation'); {\\it (ii)} that the hydrogen can be split into a diffuse intergalactic component in photoionization equilibrium and a small population of self-shielded, collisionally-ionized clumps (i.e. the highest-column-density Lyman limit systems \\cite{McQuinn_11_LyLimit}); {\\it (iii)} that non-linear corrections (including quasar duty cycles) can be ignored on sufficiently large scales \\cite{Slosar09BAO,Mesinger09_LyA_fluctuating_reionisation,White10_LyA_sims}, although I will include shot noise from the rarity of sources; {\\it (iv)} that sources averaged on large scales radiate isotropically. These seem reasonable to obtain a good estimate of the effects but in future they should be checked against numerical simulations \\updated{and more complicated analytic treatments that allow for departure from equilibrium \\cite{Zhang07_reion}}. After circulating a draft of this work, I was made aware of an independent study by Gontcho A Gontcho, Miralda-Escud\\`e and Busca (in prep); at present it seems these authors reach many similar conclusions using a different calculation framework. This is encouraging, and it will be helpful to compare our approaches in due course.  Section \\ref{sec:nonrel-radiative-transfer} develops the inhomogeneous, monochromatic radiative transfer equations; Section \\ref{sec:linearisation} discusses the application of these equations to the large-scale, linear behaviour of intergalactic \\hi. Section \\ref{sec:results} presents the main results, showing how various parameters change the distinctive imprint of radiative transfer on the intergalactic neutral hydrogen. Further discussion is given in Section \\ref{sec:discussion}, especially in relation to observations of the Lyman-$\\alpha$ forest.  Two subsidiary issues are considered in appendices. In Appendix \\ref{sec:adding-gravity}, I re-derive all equations using general relativity, so including peculiar velocities and inhomogeneous gravitational redshifting and elucidating the gauge-dependence of the results (all of which considerations turn out to impact only on scales larger than those of interest here). Appendix \\ref{sec:estimate-b_0} discusses the calibration of a particular parameter (the intergalactic \\hi bias in the absence of radiation transfer) from analytic arguments and numerical simulations.  There are a few notational matters worth settling before starting the calculation.  It is helpful to be able to decompose any quantity $X$ into its spatial mean value $X_0$ and fractional perturbations $\\delta_X$ defined by \\begin{align} \\delta_X = \\frac{X-X_0}{X_0}\\textrm{.}\\label{eq:define-delta} \\end{align} Later I will mainly deal with the Fourier transform $\\tdelta_X$ of these fractional variations; any quantity can be re-written \\begin{equation} \\delta_X(\\vec{x}) = \\frac{1}{(2\\pi)^{3/2}}\\int \\dd^3 k\\, e^{i \\vec{k} \\cdot \\vec{x}} \\tdelta_X(\\vec{k})\\textrm{,} \\label{eq:FT-delta} \\end{equation} where $\\vec{k}$ is the comoving wavevector. Finally, the power spectrum $P_X(k)$ is defined by \\begin{equation} \\langle \\tdelta_X(\\vec{k}')^* \\tdelta_X(\\vec{k}) \\rangle = P_X(k) \\delta(\\vec{k}-\\vec{k}')\\label{eq:define-Pk} \\end{equation} where angle brackets denote an ensemble average and, by an unfortunate quirk of conventional notation, the $\\delta$ on the right hand side represents the Dirac delta function. It follows from these definitions that the power spectrum for any quantity has units of a comoving volume.  The expression above assumes statistical isotropy so that $P_X$ is a function of $k=|\\vec{k}|$ alone.  Numerical results will be derived assuming a fiducial Planck temperature-only \\cite{Planck13_CosmoPar} cosmology $(h,\\Omega_{M0},\\Omega_{\\Lambda0})=(0.6711,0.3175,0.6825)$, where $\\Omega_{M0}$ and $\\Omega_{\\Lambda0}$ are the present day matter and cosmological constant densities relative to critical and $h=H_0/(100\\,\\km\\,\\mathrm{s}^{-1}\\,\\Mpc^{-1})$ is the dimensionless Hubble parameter today. The main role of these quantities will be to fix the Hubble expansion rate at $z=2.3$; any uncertainties are easily small enough to be ignored for the present study. ",
        "conclusions": "\\label{sec:discussion}  Radiative transfer imprints dramatic scale-dependent bias in the intergalactic \\hi and therefore the Lyman-$\\alpha$ forest, even after reionization is complete. This follows because regions separated by distances larger than the UV photon mean-free-path reach essentially independent photoionization equilibria. Source clustering is stronger than IGM clustering, leading to negative \\hi bias on large scales (i.e. the \\hi anti-correlates with large-scale overdensities).  This paper has presented a detailed calculation of these new effects by adopting a monochromatic, equilibrium, large-scale description, focusing on the large-scale, average correlations \\cite{Croft04_LyA_large_radiation,McDonald05,McQuinn11_LyA_fluctuations} rather than small-scale non-linear fluctuations \\cite{Viel_WDM_Lya_2008,Slosar09BAO}. The systematic analytic treatment starts from first-principles radiative transfer and produces, with minimal computational effort, predictions for very large scales.  The calculation reveals, as expected from the argument above, a strong feature in the \\hi power spectrum and a corresponding distortion of its correlation function. According to the estimates here, this distortion should have an effect at the BAO scale ($\\approx 100\\,h^{-1}\\,\\Mpc$). The BAO bump position is slightly shifted -- in Figure \\ref{fig:xi}, the local maximum is at 1.2\\% larger scales in the radiative-transfer case (solid line) compared to the constant-bias case (dashed line). That said, future cosmology constraints are unlikely to come from measuring the peak in such a simple way; so long as algorithms marginalize over possible broadband distortions to the correlation function, they will likely still recover an unbiased estimate of the BAO scale.  The most interesting conclusion is therefore that BAO-focused Lyman-$\\alpha$ observational programmes will be able to recover helpful astrophysical constraints: the radiative transfer distortions are strongly dependent on the mean bias of sources (Figures \\ref{fig:bk} and \\ref{fig:xi}), the \\hi opacity (Figure \\ref{fig:mfp}) and the effective number density of sources $\\bar{n}$ (Figure \\ref{fig:shotnoise}). As $b_j$ increases, the correlated component of the radiative fluctuations grows and the power on BAO scales decreases while the power on very large scales increases; as $\\bar{n}$ decreases, the random component of the radiative fluctuations grows and the power on all scales increases. These trends seem to agree with numerical results where a comparison can be made \\cite{Croft04_LyA_large_radiation,McDonald05,McQuinn11_LyA_fluctuations}.  The overall picture gives rise to a large number of questions. The most obvious is whether existing BOSS observations of the Lyman-$\\alpha$ forest are compatible with the expected thumbprint.  A variety of subtle observational issues must be taken into account before this can be answered. First, converting an \\hi correlation function into a flux correlation function is a non-linear process that needs at a minimum to be calibrated by suitable numerical simulations \\cite{McDonald03}. Redshift-space distortions will mix the dynamical growth of structure with the tracer statistics into a final observed correlation function \\cite{1987MNRAS.227....1K,1999coph.book.....P,McDonald03}. Dependent on the exact survey design, angular binning and data cuts, these effects could easily dilute the scale-dependence, making the observed correlation function closer to the homogeneous-radiation result. However the changes in the underlying intergalactic \\hi bias are sufficiently dramatic for it to seem implausible that the radiation-transfer signature would be obscured completely in forthcoming precision data. \\updated{To be sure we will have to understand how the data processing and parameter degeneracies impact on our ability to measure the effects. Although the distortion is large, it is also a very smooth function of scale and therefore one needs to accurately calibrate the normalization of the correlation function over a wide range of scales to make a definitive detection; otherwise the effects are degenerate with a renormalization of the uniform-limit bias $\\bu$.}  The BOSS team have emphasized that their 3D Lyman-$\\alpha$ forest pipeline is presently designed to pick out localized correlation-function features -- i.e. the BAO bump -- rather than reconstruct the entire function accurately \\cite{2013_Busca_BOSS_Lya_BAO,BOSS_Slosar13}. Nonetheless an attempt to measure scale-dependence in cross-correlation against quasars has been made; none was found \\cite{FontRibera_BOSS_QSO_Lya_X_2013}. Conversely some scale-dependence in the cross-correlation between damped Lyman alpha systems and the forest can be seen in the plots of Ref.  \\cite{2012FontRibera_BOSS_DLA_Forest_X}.  It is unclear whether and how these results can be reconciled with the present work; observational difficulties such as continuum determination cause severely correlated errors in correlation functions and dealing with these leads to certain large-scale modes being projected out \\cite{2013_Busca_BOSS_Lya_BAO,BOSS_Slosar13}. Overall, the task of determining whether the effects of radiative transfer are present in existing data is considerable. However, I hope that the present calculation has underlined the rewards of such an effort. The scale-dependent radiative transfer contains a rich, valuable source of information on the nature of UV sources.  21cm emission studies will not be affected by these considerations because the \\hi they probe is largely in collisional- rather than photo-ionization equilibrium \\cite{PontzenDLA}. The 21cm absorption forest would be affected in just the same way as the Lyman-$\\alpha$ forest; but this phenomena is of most promise at high redshift before or during reionization \\cite{Pritchard21cm_review} -- so the present calculation does not apply. \\updated{One way to tackle the larger fluctuations at high redshift is to use a halo-model-based calculation \\cite{Furlanetto04_reion_model}; alternatively, a linear theory approach has been taken to the problem by Ref. \\cite{Zhang07_reion,Aloisio13_reion_fNL,Mao13_reion_fNL}. In these cases, an explicit time integration needs to be performed to follow the growth of ionized bubbles whereas in the present case the integration is absent because I have assumed equilibrium, making the present paper's calculations considerably simpler but more restricted in scope.  }  Depending on one's assumptions (for instance on the relative importance of quasars to the UV background, and on the quasar luminosity function), the rarity of sources also have a substantial impact on very large scales. Here I have modeled the resulting shot noise by a Gaussian approximation similar to that of Ref. \\cite{McQuinn11_LyA_fluctuations}; \\updated{in that work, noise was considered to be so large that the correlated component of the radiation fluctuations was thrown out of the calculation. In the present work the effects of shot-noise seem milder, which reflects that I work at lower redshift (where the comoving density of quasars has increased) and make greater allowance for a UV contribution from star-forming galaxies.} Crucially, the correlated component has a qualitatively different signature to the noise component of the radiation field: the former reduces the power of \\hi fluctuations on large scales, whereas the latter can only ever add power (at any scale). One effect that is absent from the present work concerns scales below $10\\,\\Mpc$ or so -- here the noise would be significantly amplified \\cite{Mesinger09_LyA_fluctuating_reionisation} by 2-halo and other nonlinear effects. Another missing aspect from my analysis is that of time-dependence which could, for example, add further confusion from quasar duty cycles.  With all this in mind it would be of great interest to supplement the linear theory calculations of this work by revisiting the BAO-scale correlation function of the Lyman-$\\alpha$ forest using non-linear numerical simulations of gigaparsec chunks of the IGM with correlated sources, incorporating radiative transfer -- along the lines of work described by Refs \\cite{Croft04_LyA_large_radiation,McDonald05,McQuinn11_LyA_fluctuations}. Hints of the anti-correlation discussed at length in the present paper have been seen before in such efforts \\cite{Croft04_LyA_large_radiation,White10_LyA_sims}. It would be helpful to include large scale temperature fluctuations arising from helium reionization \\cite{Furlanetto09_He_fluc}. Or, one might be able to tackle temperature fluctuations analytically by relaxing the monochromatic assumption; it is worth re-emphasizing that the current work includes the zero-order effects of hard photons (the spectral shape enters through equation \\eqref{eq:alphaz}), but not first-order changes from local fluctuations in spectral shape. \\updated{At this level of approximation, the gas thermal equilibrium is nearly unaffected by radiation intensity fluctuations \\cite{HuiGnedin97_IGM_EoS} because the radiative heating rate can be approximately re-written as a function of density and temperature (via the ionization equilibrium condition). To answer the important question of how thermal fluctuations change the large-scale signal one therefore needs either to incorporate multi-wavelength radiative transfer or go beyond an equilibrium approximation -- or, preferably, both \\cite{HuiGnedin97_IGM_EoS}.}  Further work is required to reach a unified view of how radiation changes the observational prospects for cosmology and astrophysics with the Lyman-$\\alpha$ forest. The present investigation forms a first guide to the effects that will dominate on the largest scales."
    },
    "1402/1402.5053_arXiv.txt": {
        "abstract": "{The object HD\\,43587Aa is a G0V star observed during the 145-day LRa03 run of the COnvection, ROtation and planetary Transits space mission (CoRoT), for which complementary High Accuracy Radial velocity Planet Searcher (HARPS) spectra with S/N$>$300 were also obtained. Its visual magnitude is 5.71, and its effective temperature is close to 5950\\,K. It has a known companion in a highly eccentric orbit and is also coupled with two more distant companions.}{We undertake a preliminary investigation of the internal structure of HD\\,43587Aa.}{We carried out a seismic analysis of the star, using maximum likelihood estimators and Markov Chain Monte Carlo methods.} {We established the first table of the eigenmode frequencies, widths, and heights for HD\\,43587Aa. The star appears to have a mass and a radius slightly larger than the Sun, and is slightly older ($5.6$ Gyr). Two scenarios are suggested for the geometry of the star: either its inclination angle is very low, or the rotation velocity of the star is very low.} {A more detailed study of the rotation and of the magnetic and chromospheric activity for this star is needed, and will be the subject of a further study. New high resolution spectrometric observations should be performed for at least several months in duration.} ",
        "introduction": "\\label{intro} The star HD\\,43587Aa is a primary target of the asteroseismology program of the COnvection, ROtation and planetary Transits space mission (CoRoT)\\footnote{The CoRoT space mission, launched on 2006 December 27, was developed and is operated by the Centre national d'\u00e9tudes spatiales (CNES) with participation of the Science Programs of the European Space Agency (ESA), ESA's Research and Scientific Support Department (RSSD), Austria, Belgium, Brazil, Germany and Spain.}. CoRoT observed ten main sequence stars displaying solar-like oscillations \\citep{Michel08}. This set of pulsators consists of: five F-stars, HD\\,49933, HD\\,181420, HD\\,181906, HD\\,175726, and HD\\,170987 (see \\citealt{Appourchaux08,Benomar2009b,Barban09,Garcia_2009,Mosser_2009,Mathur_2010}); the evolved G star HD\\,49385  \\citep{Deheuvels10}; the K and G stars HD\\,46375 and HD\\,52265, which host non-transiting planets (\\citealt{Gaulme_2010,Ballot11}); the G star HD\\,169392, which belongs to a weakly bound binary system \\citep{Mathur_2013}; and the star that is the subject of the present study HD\\, 43587Aa (G0V).\\  Binary systems are fundamental astrophysical objects: when coupled with asteroseismology, they may provide two independent methods to obtain masses and radii, and therefore offer opportunities to test internal structure modelling and develop highly constrained stellar models. Indeed, it is occasionally possible to determine the masses and/or radii of stars constituting a binary or a multiple star system, provided that the orbital parameters are measurable. These can be obtained either by optically observing the components of visual binaries, or by combining photometric and radial velocity measurements in the case of eclipsing binaries. An alternative way of retrieving masses in photometry, for close-in binary systems, consists of modelling the photometric modulation resulting from the tidally distorted shapes of the stars, which are a function of stellar masses (e.g., \\citealt{Mazeh_2008}). Finally, masses of triple system components, where at least two components mutually eclipse, can be deduced from eclipse timing variations (e.g., \\citealt{Steffen_2013}). Recently, efforts were made to exploit the combination of binarity and asteroseismology, in particular with data from the NASA \\textit{Kepler} mission. \\citet{Hekker_2010} and \\citet{Frand13} reported the identification of the eclipsing binary system KIC\\,8410637, composed of an F star and a red giant that exhibit solar-like oscillations. \\citet{Metcalfe_2012} performed the seismic analysis of the binary system 16 Cyg A and B, where eigenmodes of oscillations were identified for both solar-type stars. More recently, \\citet{Gaulme_2013} reported the detection of 12 new eclipsing binary systems, one non-eclipsing binary with tidally induced oscillations, and ten more candidate triple systems, all of which include a pulsating red giant.\\    \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{HD43587_figure1_rev}} \\caption{Light curve of HD\\,43587 (LRa03 run of CoRoT), after a second-order polynomial fit was removed. A close-up of 80 minutes displays the temporal signature of the most powerful eigenmodes.} \\label{lightcurve} \\end{figure}    The subject of our work, HD\\,43587Aa, a G0V star whose visual apparent magnitude is $m_V=5.71$, belongs to a quadruple system composed of two distant main sequence visual binaries (HD\\,43587~A \\& E) as imaged by \\citet{Hart12} and by \\citet{Pravdo06}, respectively. The denomination A and E is that of the Washington Double Star catalogue (WDS). We were able to check that HD\\,43587A and HD\\,43587E had common proper motions and parallaxes measured with sufficient precision \\citep{Lepi05}. Thanks to radial velocity measurements, \\citet{Vogt02} were the first to determine estimates of the orbital properties of HD\\,43587A, followed by \\citet{Catala06} who used adaptive optics to resolve both stars in the system. Recently, \\citet{Katoh13} found the most precise values to date using high-dispersion spectroscopy. The orbital period, eccentricity, and semi-major axis are estimated to be $P = 34.2 \\pm 0.2$~ years, $e = 0.8045 \\pm 0.0009$, and $a = 11.61 \\pm 0.11$\\,AU respectively. \\citet{Katoh13} gave also estimates of $1.01\\pm 0.02~M_\\odot$ and $5926 \\pm ~80$ K for the mass and the effective temperature of  HD\\,43587Aa; they also derived a minimum value of $0.342\\pm 0.003~M_\\odot$ for the close companion HD\\,43587Ab. The system HD\\,43587A thus appears to be composed of a solar analog and an M star (effective temperature of $3820 \\pm 100$ K following \\citealt{Catala06}), and to orbit in a highly eccentric orbit.\\\\  A spectroscopic analysis was carried out by \\citet{Bruntt04} who computed effective temperatures via different methods: 5923\\,K (line depth ratio), 5850\\,K (H$\\alpha$ wings), and 5931\\,K (Stromgren indices). Recently, \\citet{Morel13} published a rather precise value, based on new data: $T_{\\rm eff} = 5947\\pm17$\\,K  (more details are given in section 2.1). \\citet{Vogt02} measured $v\\,\\sin i$ to be 2.7\\,km/s, while \\citet{Bruntt04} suggested a value of 2.5\\,km/s even though this is an extrapolation beyond the resolution limit of the spectrograph they used. The star has a revised Hipparcos parallax of 52.0\\,mas and a metallicity of $-0.11$ \\citep[in the G-C catalogue and][]{Bruntt04}. \\citet{Bruntt04} estimated the surface gravity to be 4.31 from Stromgren indices, and 4.29 from evolutionary tracks and parallax. As far as the radius is concerned, a careful calculation was done by \\citet{Thevenin06} using a calibrated Barnes-Evans relation to obtain a limb-darkening estimate and they found a value of $R = 1.280\\pm0.032 R_{\\odot}$.\\  We will present in Section 2 the spectrometric ground-based support that was specifically performed to refine the temperature and abundance estimates, as well as the CoRoT photometric data that led to the identification of a high signal-to-noise ratio oscillation spectrum. Next, we describe how we measure the modes' properties (Sect. 3) and interpret them in terms of stellar bulk properties (Sect. 4). Finally, the modelling of the stars' internal structure with an evolutionary code allows us to estimate its parameters to an unprecedented precision (Sect. 5). For the sake of simplicity, HD\\,43587Aa will be called HD\\,43587 hereafter.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{HD43587_figure2.eps}} \\caption{Power spectrum of HD\\,43587 in the p-mode range.} \\label{dse} \\end{figure} ",
        "conclusions": "The analysis of the light curve obtained by CoRoT of HD\\,43587A exhibits clear acoustic modes, which allowed us to establish a reference frequency table. Time-frequency analysis and MCMC methods were used to help discriminate between noise and acoustic signal when the signal-to-noise ratio was low. The detection of several  $\\ell$\\,=\\,3 modes  proved to be statistically significant. The frequency splitting induced by the internal rotation of the star could not be extracted satisfactorily, its probability density function being bimodal. Two scenarios could be realistic for the geometry of the star: either its inclination angle is very low, which disminishes the signature of rotation in the spectrum, or the rotation of the star is very low, with the same consequence on the spectrum. No conclusion could be reached concerning the surface rotation itself, with no magnetic structures due to the extremely low level of activity.  The HD\\,43587Aa seismic model corresponds to a star slightly more massive, hotter, and thus more luminous than the Sun: $M_\\mathrm{model} = 1.04 M_{\\odot};\\, R_\\mathrm{model} = 1.19R_{\\odot};\\,L_\\mathrm{model}=1.58\\,L_{\\odot}$ and an age of $5.6$ Gyr. The rotation profile remains particularly challenging: a solar analogue with a very slow rotation would result from a different history from that of the Sun. The physics that could be responsible for this (protostellar cloud, early age of the star, composition, magnetic field) remains uncertain, as well as the impact of the binarity of the HD\\,43587A system on the evolution and the properties of HD\\,43587Aa. New high resolution spectrometric observations should be realised with the sampling and the duration required to be able to decrease the upper limit of the rotation rate. A more precise analysis of the magnetic and chromospheric activity for this star is needed, and will be the subject of a further study."
    },
    "1402/1402.5573_arXiv.txt": {
        "abstract": "The target asteroid of the OSIRIS-REx asteroid sample return mission, (101955) Bennu (formerly 1999~RQ$_{36}$), is a half-kilometer near-Earth asteroid with an extraordinarily well constrained orbit. An extensive data set of optical astrometry from 1999--2013 and high-quality radar delay measurements to Bennu in 1999, 2005, and 2011 reveal the action of the Yarkovsky effect, with a mean semimajor axis drift rate $da/dt = (-19.0 \\pm 0.1)\\times 10^{-4}$ au/Myr or $284\\pm 1.5\\;\\rm{m/yr}$. The accuracy of this result depends critically on the fidelity of the observational and dynamical model. As an example, neglecting the relativistic perturbations of the Earth during close approaches affects the orbit with $3\\sigma$ significance in $da/dt$.  The orbital deviations from purely gravitational dynamics allow us to deduce the acceleration of the Yarkovsky effect, while the known physical characterization of Bennu allows us to independently model the force due to thermal emissions. The combination of these two analyses yields a bulk density of $\\rho = 1260\\pm70\\,\\rm{kg/m^3}$, which indicates a macroporosity in the range $40\\pm10$\\% for the bulk densities of likely analog meteorites, suggesting a rubble-pile internal structure. The associated mass estimate is $(7.8\\pm0.9)\\times 10^{10}\\, \\rm{kg}$ and $GM = 5.2\\pm0.6\\,\\rm{m^3/s^2}$.  Bennu's Earth close approaches are deterministic over the interval 1654--2135, beyond which the predictions are statistical in nature. In particular, the 2135 close approach is likely within the lunar distance and leads to strong scattering and therefore numerous potential impacts in subsequent years, from 2175--2196. The highest individual impact probability is $9.5\\times 10^{-5}$ in 2196, and the cumulative impact probability is $3.7\\times 10^{-4}$, leading to a cumulative Palermo Scale of -1.70. ",
        "introduction": "\\label{sec:intro}  The Apollo asteroid (101955) Bennu, a half-kilometer near-Earth asteroid previously designated 1999~RQ$_{36}$, is the target of the OSIRIS-REx sample return mission. A prime objective of the mission is to measure the Yarkovsky effect on this asteroid and constrain the properties that contribute to this effect. This objective is satisfied both by direct measurement of the acceleration imparted by anisotropic emission of thermal radiation, the first results of which are reported here, and by constructing a global thermophysical model of the asteroid to confirm the underlying principles that give rise to this effect.  Bennu was discovered by the LINEAR asteroid survey in September 1999. Since then, more than 500 optical observations have been obtained for this Potentially Hazardous Asteroid (PHA).  Moreover, the asteroid was observed using radar by the Arecibo and Goldstone radio telescopes during three different apparitions. Thanks to this rich observational data set, Bennu has one of the most precise orbits in the catalog of known near-Earth asteroids. The exceptional precision of the Bennu orbit allows one to push the horizon for predicting possible Earth impacts beyond the 100 years typically used for impact monitoring \\citep{milani_etal_nonlinear}, and indeed \\citet{milani_etal_rq36} showed that Earth impacts for Bennu are possible in the second half of the next century. In particular, the cumulative impact probability they found was approximately $10^{-3}$, about half of which was associated with a possible impact in 2182. However, the occurrence of an impact depends decisively on the Yarkovsky effect because the prediction uncertainty due to this nongravitational perturbations dominates over the orbital uncertainty associated with astrometric errors.  The Yarkovsky effect is a subtle nongravitational perturbation that primarily acts as a secular variation in semimajor axis and thus causes a runoff in orbital anomaly that accumulates quadratically with time \\citep{bottke_etal06}. The computation of the Yarkovsky perturbation requires a rather complete physical model of the asteroid, including size, shape, density, spin rate and orientation, thermal properties, and even surface roughness \\citep{rozitis_green_yarko}. Though such a complete profile is rarely available, the orbital drift due to the Yarkovsky effect can sometimes be determined from an asteroid observational data set. For example, \\citet{chesley_golevka} managed to directly estimate the Yarkovsky effect for asteroid (6489) Golevka by using three radar apparitions. \\citet{1992bf} employ the Yarkovsky effect to match precovery observations of asteroid (152563) 1992~BF that are incompatible with purely gravitational dynamics. More recently \\citet{nugent_etal_2012} and \\citet{farnocchia_yarkolist_2013} have estimated the Yarkovsky effect for a few tens of near-Earth asteroids by using a formulation that depends on a single parameter to be determined from the orbital fit.  Besides Bennu, there are two other asteroids for which possible impacts are known to be driven by the Yarkovsky effect: (29075) 1950~DA \\citep{giorgini_1950da} and (99942) Apophis \\citep{chesley_acm05, giorgini_apophis_08}. The relevance of the Yarkovsky effect for Apophis is due to a scattering close approach in 2029 with minimum geocentric distance $\\sim$38000 km. For 1950 DA the influence of the Yarkovsky effect for an impact in 2880 is due to the long time interval preceding the potential impact. However, no estimate of the Yarkovsky perturbation acting on these two asteroids is currently available. To analyze such cases one can use the available physical constraints for the specific objects, along with general properties of near-Earth asteroids (e.g., albedo, thermal inertia, bulk density, etc.) to statistically model the Yarkovsky effect. The orbital predictions and the impact hazard assessment are then performed by a Monte Carlo simulation that accounts for both the Yarkovsky effect distribution and the orbital uncertainty \\citep{farnocchia_etal_apophis, farnocchia_etal_1950da}. For Bennu, no such heroics are required. As we shall see, we now have a precise estimate of the orbital deviations caused by the Yarkovsky effect, as well as a comprehensive physical model distilled from numerous investigations.  While the Yarkovsky effect requires a priori knowledge of several physical parameters to be computed directly, its detection through orbital deviations can be used to constrain the otherwise unknown physical parameters. When the spin state is unknown, one can derive weak constraints on obliquity, as was first shown by \\citet{1992bf} for 1992 BF. In cases where the spin state is well characterized, usually through the combination of radar imaging and photometric light curves, the bulk density of the object is correlated with the thermal properties and mutual constraints can be inferred, as was the case for Golevka \\citep{chesley_golevka}. \\citet{rozitis_2013_apollo} were able to jointly model the measured Yarkovsky and YORP effects on (1862) Apollo, and thereby constrain a number of the body's physical characteristics, including axis ratios, size, albedo, thermal inertia and bulk density. In the case of Bennu, the thermal inertia is known from infrared observations \\citep{emery_bennu_2014, muller_2012_rq36}, and so we are able to directly estimate the mass and bulk density. ",
        "conclusions": "Understanding of an asteroid's physical properties becomes essential whenever the Yarkovsky effect or other nongravitational accelerations are a crucial aspect of the orbit estimation problem. Radar astrometry of asteroids can provide surprising and important constraints, not only on an asteroid's orbit, but also on its physical properties. In the case of Bennu, this information has immense value for space mission designers. We have seen that the availability of well-distributed radar astrometry over time spans of order  a decade can constrain asteroid orbits to the extent that precise estimates of the Yarkovsky effect can be derived. When coupled with thermal inertia information derived from other sources, such as the Spitzer Space Telescope, important parameters such as mass, bulk density and porosity can be derived. Combining Yarkovsky detections with thermal inertia measurements to infer the asteroid mass can be implemented on other near-Earth asteroids, including potential space mission targets. This technique is the focus of ongoing work. Indeed, Bennu clearly demonstrates that even weak radar detections can have considerable science value, raising the imperative to aggressively pursue every available radar ranging opportunity for potential Yarkovsky candidates.  Our bulk density estimate for Bennu implies a primitive body with high porosity of $40\\pm10$\\%. The implication is that Bennu must be comprised of a gravitationally bound aggregate of rubble, a conclusion that is reinforced by its shape, which is spheroidal with an equatorial bulge consistent with the downslope movement and accumulation of loose material at the potential minimum found at the equator \\citep{nolan_etal_2013}. This bodes well for the OSIRIS-REx sample collection effort, which requires loose material at the surface for a successful sample collection, although nothing in this study constrains the size distribution of the surface material.  The statistical encounter frequency with Earth (Fig.~\\ref{fig:prob_ca}) can be used to understand the rate of encounters that could alter the shape and spin state of a body through tidal interactions \\citep[e.g.,][]{walsh_richardson_2008, nesvorny_etal_2010}. \\citet{scheeres_apophis_2005} have shown that tidal interactions at a distance of ~6 Earth radii can appreciably alter the spin state of 99942 Apophis. More distant encounters could still excite the spin state enough to induce seismic activity, leading to a periodic resurfacing of the asteroid that may have implications for interpretation of Bennu samples returned by OSIRIS-REx.  We have seen that the current levels of uncertainty in Bennu's orbit are low enough that unprecedented levels of accuracy are required in the dynamical model that governs the trajectory. For example, the relativistic perturbation of planetary gravity fields, in particular that of Earth, must be incorporated to obtain reliable results. The future addition of OSIRIS-REx radio science data will again decrease the orbital uncertainties by 1--2 orders of magnitude, which will likely require even finer scale refinements to our dynamical model than used here. However, difficulties in understanding the proper statistical treatment of asteroid optical astrometry, and in particular the identification of statistical outliers, will likely remain a dominant source of uncertainty that is not properly captured by a posteriori covariance analysis.  Thus our findings for the post-OSIRIS-REx orbital uncertainties of Bennu may be illusory. The finding that $\\zeta_{2135}$ uncertainties may be reduced as low as 1000 km with the OSIRIS-REx radio science data assumes that our Yarkovsky model, including the asteroid spin state, holds through 2135. Thus a host of model refinements may be necessary to properly characterize the trajectory out to 2135. Notwithstanding the next radar observation opportunity in January 2037, we may reach an uncertainty limit that prevents us from improving predictions any further until models can improve or the prediction interval is significantly reduced. As an example, the post-2135 predictability will markedly improve after the 0.005 au Earth close approach in 2060, and it is reasonable to expect that at least some potential impacts will persist until that time."
    },
    "1402/1402.2353_arXiv.txt": {
        "abstract": "{A patch of sky in the SDSS Stripe 82 was observed at 1.6~GHz with Very Long Baseline Interferometry (VLBI) using the European VLBI Network (EVN). The data were correlated at the EVN software correlator at JIVE (SFXC). There are fifteen known mJy/sub-mJy radio sources in the target field defined by the primary beam size of a typical 30-m class EVN radio telescope. The source of particular interest is a recently identified high-redshift radio quasar J222843.54+011032.2 (J2228+0110) at redshift $z$=5.95. Our aim was to investigate the milli-arcsecond (mas) scale properties of all the VLBI-detectable sources within this primary beam area with  a diameter of 20$\\arcmin$. The source J2228+0110 was detected with VLBI with a brightness temperature $T_{\\rm b}$$>$$10^8$~K, supporting the active galactic nucleus (AGN) origin of its radio emission, which is conclusive evidence that the source is a radio quasar. In addition, two other target sources were also detected, one of them with no redshift information. Their brightness temperature values ($T_{\\rm b}$$>$$10^7$~K) measured with VLBI suggest a non-thermal synchrotron radiation origin for their radio emission. The detection rate of 20\\% is broadly consistent with other wide-field VLBI experiments carried out recently. We also derived the accurate equatorial coordinates of the three detected sources using the phase-referencing technique. This experiment is an early attempt of a wide-field science project with SFXC, paving the way for the EVN to conduct a large-scale VLBI survey in the multiple-phase-centre mode.} ",
        "introduction": "Very Long Baseline Interferometry (VLBI) surveys of extragalactic radio sources supply a wealth of information on the nature of physical phenomena in the close vicinity of active galactic nuclei (AGNs), with both astrophysical and cosmological applications \\citep[for reviews, see e.g.][]{Gur04,Fre06}. However, traditional VLBI observations have very small fields of view compared to the average angular distance between compact extragalactic radio sources at the achievable sensitivity. Therefore, until recently, VLBI surveys are usually restricted to targeted observations of bright sources, and thus the growth of VLBI-detected AGN samples has been relatively slow over the past decades.  The high data recording rate of modern VLBI and long coherent on-source integration times achieved by the phase-referencing technique allow us to detect mJy-level or even sub-mJy-level weak sources. In its traditional way, phase-referencing involves regularly interleaving observations between the target source and a nearby, bright, and compact phase-reference calibrator \\citep{Bea95}. The participating telescopes in a VLBI network change their pointing direction, nodding between the target and the reference source, with a period shorter than the atmospheric coherence scale. While VLBI surveys exploiting this technique are technically feasible \\citep[e.g.][]{Gar99,Mos06,Fre08a}, they require a prohibitively large amount of VLBI observing resources to obtain deep mas-resolution radio imaging information on thousands of faint AGNs.  With the technique of in-beam phase-referencing \\citep[e.g.][]{Gar01,Gar05,Len08}, any additional source that is present in the primary beam of the individual radio telescopes of the interferometer network can be phase-calibrated with a bright and compact primary object, without repointing the antennas. However, to observe a field as extended as possible by minimizing bandwidth and time-average smearing, experiments with a single correlation pass would require extremely high output data rates and thus produce very large data files, which are computationally challenging to analyse \\citep[e.g.][]{Chi13}. Another possibility is to use as many different correlator passes as the number of targets, but then the required correlator resources are very considerable if not prohibitive.  The advent of multiple-phase-centre processing \\citep{Del11,Mor11} using software VLBI correlators \\citep[e.g.][]{Del07,Del11,Pid09} offers a possibility to produce datasets with many different phase centres located within the primary beam in a single correlator pass. The individual output file for each target has the usual averaging in time and frequency, and is easily manageable in the process of the calibration and imaging. Thus the in-beam phase referencing combined with the multiple-phase-centre correlation dramatically decreases the demand on the observing resources, as well as on computational requirements, while it only moderately increases the required correlator time. This in fact multiplies the surveying power of the VLBI network by a substantial factor. It has become feasible to plan and conduct systematic surveys of thousands of sources around suitable in-beam calibrators \\citep{Del14}. This breakthrough will allow us to reveal mas-scale radio structures in a large sample of extragalactic sources, while pushing down the flux density limit by at least two orders of magnitude. In turn, this will bring an order of magnitude more VLBI-imaged extragalactic sources than are accessible for studies at present.  It is hard to predict all the novel science applications that will become possible with significantly (orders of magnitude) increased size samples of VLBI-imaged extragalactic radio sources. Relevant research fields include studies of the cosmological evolution of active galaxy populations, the role of star formation and accretion onto the central supermassive black hole, and their interplay (feedback mechanisms). Extensive optical sky surveys, such as the Sloan Digital Sky Survey \\citep[SDSS,][]{Yor00} and the prospective surveys by the Large Synoptic Survey Telescope \\citep[LSST,][]{Ive08}, reduce the problem of missing optical identifications and spectroscopic redshifts. High-resolution radio surveys can be matched with large multi-wavelength databases, such as in the infrared and X-ray bands, to study dust obscuration in the AGN host galaxies. Morphological classification of low-luminosity sources could eventually be made using a large sample. The properties of these sources can then be compared with those of well-studied powerful objects. Thousands of new radio AGNs imaged with VLBI and having measured cosmological redshift could provide the necessary impetus for renewed attempts to estimate cosmological model parameters using ``standard'' objects. Finally, collecting more AGNs with radio and optical identification could lead to a more accurate and dense astrometric link of the radio and optical reference frames.  At the 1.4-GHz flux densities $S_{1.4}$$>$1~mJy, the radio source population is believed to be dominated by radio-loud AGNs, and the star-forming galaxies start to become a major component of the sub-mJy population \\citep[e.g.][]{Nor11}. Nevertheless, the exact fraction of AGNs among the sub-mJy radio sources is unknown. Using the multiple-phase-centre technique, \\citet*{Mid11} observed an area containing the Chandra Deep Field South (CDFS) with the Very Long Baseline Array (VLBA), obtained a detection rate of 21\\%, and identified seven new AGNs. The work, together with other wide-field experiments carried out recently \\citep[e.g.][]{Mid13,Mor13}, demonstrated the capability of VLBI as a powerful tool for unveiling the nature of mJy/sub-mJy radio sources in multiple-phase-centre mode. In particular, this method can help identify the heavily obscured AGNs via their the radio emission that is completely free of dust obscuration.  The highest redshift ($z$$>$5.7) radio quasars belong to the mJy/sub-mJy radio source population due to their large luminosity distance. So far, three mJy-level $z$$\\sim$6 quasars have been observed with VLBI \\citep{Fre03,Fre05,Mom08,Fre08b,Fre11}. Their steep radio continuum spectra (spectral index $\\alpha$$<$$-0.5$; $S_{\\nu} \\varpropto \\nu^{\\alpha}$, where $S_{\\nu}$ is the flux density and $\\nu$ the observing frequency) and compact radio structures ($<$$100$~pc) show that they have characteristics similar to the gigahertz-peaked-spectrum (GPS) sources known at lower redshifts. However, a larger sample is needed to draw a firm conclusion about whether most (if not all) of the highest redshift radio quasars are truly young radio sources in the early Universe.  The GPS and CSS (compact-steep-spectrum) sources are dominated by somewhat extended emission, although they both are confined to compact regions. The GPS sources have a size smaller than that of the narrow-line region (i.e. $<$1 kpc), while the CSS sources are entirely contained in their host galaxies (i.e. $<$15 kpc). They have a convex radio spectrum and barely show any variability in luminosity. They may be newly born radio sources that will eventually evolve into large Fanaroff--Riley \\citep{Fan74} Class I (FR~I) or Class II (FR~II) radio sources (i.e. the youth scenario) or that are frustrated by the ambient dense gas/dust clouds \\citep[i.e. the frustration scenario; see][for a review]{Dea98}. For the highest redshift GPS/CSS sources, both scenarios could be present. The very first generation of supermassive black holes had just recently developed their jets and expanding double mini-lobes, which at the same time were confined by the extremely dense interstellar medium (ISM) of their forming host galaxies \\citep{Fal04}.  Recently, \\citet*{Zei11} have identified the second radio-selected $z$$\\sim$6 quasar \\object{J222843.54+011032.2} (J2228+0110, hereafter) by matching the optical detections of the deep SDSS Stripe 82 and their radio counterparts in the Stripe 82 Very Large Array (VLA) Survey{\\footnote{\\url{http://sundog.stsci.edu/cgi-bin/searchstripe82}} \\citep{Hod11}. The integrated flux density of $S_{1.4}$=1.32~mJy means that the source is the fourth $z$$\\sim$6 radio-luminous quasar known to date with $S_{1.4}$$>$ 1~mJy and therefore is a feasible target for VLBI to probe the physical properties of the rare radio sources at the highest redshifts ($z$$\\sim$6) at milli-arcsecond (mas) scales using the phase-referencing technique. In the project described in this paper, we observed J2228+0110 with the European VLBI Network (EVN) in the multiple-phase-centre mode, to simultaneously explore the mas-scale features of the neighbouring mJy/sub-mJy radio sources falling in the field of view (FOV).  In this paper, we adopt a flat $\\Lambda$CDM cosmological model, with $H_0$=71~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}$=0.27, and $\\Omega_\\Lambda$=0.73, which yields a projected linear scale of 5.86~pc\\,mas$^{-1}$ at $z$=5.95. Our selection method for the field and the targets is described in Section~\\ref{field}. The observations and the data reduction procedures are presented in Section~\\ref{observations}. The results are shown in Section~\\ref{results} and discussed in Section~\\ref{discussion}. Some conclusions are drawn in Section~\\ref{conclusions}.  \\begin{table} \\caption{Fifteen target sources in the EVN experiment described in this paper.} \\begin{center} \\scalebox{0.9}{ \\begin{tabular}{ccccc} \\hline \\hline Source name & RA (J2000) & DEC (J2000) &  Flux density \\\\ & h m s      & $\\degr$ $\\arcmin$ $\\arcsec$ &  mJy   \\\\ \\hline FIRST-6     &  22 28 42.206 &  01 14 47.73 &  1.06   \\\\ J2228+0110\\,* &  22 28 43.540 &  01 10 32.22 &  1.32  \\\\ VLA-5\\,*    &  22 28 49.842 &  01 13 37.47 &  1.76  \\\\ NVSS-1      &  22 28 51.380 &  01 07 09.40 &  2.5    \\\\ VLA-4\\,*    &  22 28 51.452 &  01 12 03.39 &  0.37  \\\\ NVSS-2      &  22 29 00.420 &  01 17 46.90 &  2.7    \\\\ VLA-3\\,*    &  22 29 18.059 &  01 06 18.43 &  7.97  \\\\ VLA-1\\,*    &  22 29 18.689 &  01 09 03.09 &  1.44  \\\\ VLA-2\\,*    &  22 29 25.487 &  01 17 52.61 &  0.43  \\\\ FIRST-1     &  22 29 31.387 &  01 14 43.99 &  0.90   \\\\ FIRST-2     &  22 29 33.802 &  01 15 38.08 &  2.02   \\\\ FIRST-4     &  22 29 38.594 &  01 17 34.96 &  1.02   \\\\ NVSS-3      &  22 29 38.790 &  01 18 05.20 &  3.1    \\\\ FIRST-3     &  22 29 41.761 &  01 14 28.46 &  2.07   \\\\ FIRST-5     &  22 29 52.917 &  01 12 19.57 &  0.92   \\\\ \\hline \\end{tabular}} \\end{center} \\tablefoot{Col.~2--3: a priori right ascension (RA) and declination (DEC) in the J2000 equatorial coordinate system, which were used for the correlation; the coordinates of the sources marked with an asterisk are from the Stripe 82 VLA Survey \\citep{Hod11}; the rest are from the FIRST \\citep{Whi97} and NVSS \\citep{Con98} surveys, as indicated by the source names; Col.~4: the integrated flux density of each source at 1.4~GHz measured in the respective low-resolution VLA observations mentioned above. \\label{tbl-1}} \\end{table} ",
        "conclusions": "\\label{conclusions} We observed and detected the second radio-selected $z$$\\sim$6 quasar J2228+0110 with the EVN in wide-field mode, using in-beam phase referencing. In addition, with the same observing data, we applied the novel multiple-phase-centre correlation technique and detected two other target radio sources within a 20$\\arcmin$-diameter field. The measured brightness temperature of J2228+0110 is similar to that of the three other known $z$$>$5.7 radio quasars, and the inferred size suggests it might be a CSS source. The high brightness temperature of another detected object, VLA-4 (\\object{J222851.44+011203.4}), together with the photometric redshift of 0.74$\\pm$0.13, implies AGN activity in its host galaxy. The nature of FIRST-3 (\\object{J222941.76+011427.4}) remains unclear in the absence of optical identification and redshift information. It is most likely an obscured AGN. The accurate astrometic coordinates of the three targets were also derived using the phase-referencing technique. The detection rate of 20\\% is broadly consistent with the findings of larger surveys, but is lower than expected for the sub-sample with $S_{1.4}$$>$1~mJy. The result confirms that an important fraction of mJy/sub-mJy radio sources are (at least partially) powered by AGN.  With the project presented here, and while successfully imaging a rare high-redshift radio quasar, we demonstrated the ability of the EVN to obtain high-resolution structural information on unrelated radio sources that happen to lie within $\\sim$10$\\arcmin$ of the primary phase centre. Experiments like this could later be used for exploring the close vicinity of any strong calibrator source observed with the VLBI network for sufficiently long time, irrespective of the main science goal. This approach does not place any special requirement on VLBI array pointing and scheduling and, in particular, does not require any excess observing time from the VLBI network. Therefore the data can be obtained very efficiently, in piggyback mode. The only extra effort needed is the multiple-phase-centre correlation of the recorded data. In this paper, we demonstrated how to analyse the data from such an experiment."
    },
    "1402/1402.4605_arXiv.txt": {
        "abstract": "Gamma Ray Bursts (GRBs) are found at much higher redshifts (z$>$6) than Supernova Ia (z$\\sim$1), and hence, they can be used to probe very primitive universe. However, radiation mechanism of GRB remains a puzzle, unlike Supernova Ia. Through comprehensive description, both empirical and physical, we shall discuss the most likely way to use the constituent pulses of a GRB to find the radiation mechanism as well as using the pulses as luminosity indicators. ",
        "introduction": "First discovered during late 1960's, Gamma Ray Burst (GRB) soon became one of the greatest puzzles in astrophysics in terms of its location, size and energetics. Many satellites have been flown since then to understand this puzzling phenomenon --- \\emph{HETE-2}, \\emph{BATSE}, \\emph{Swift}, \\emph{Fermi}, to name a few. Though we have a clear idea about their cosmic origin (the highest spectroscopic redshift being 8.2 for GRB 090423) and a rough idea about their energy budget (highest among all astrophysical phenomena --- $\\sim 10^{52}$ erg), the emission mechanism is still unsettled. Hence, in spite of a great hope that GRB could be used as luminosity indicator in extension to the currently used ones, e.g., supernova Ia, one has to standardize the GRB energetics first. This, of course, serves two purposes --- solving the GRB physics, which is not settled till now, and using GRB as luminosity indicator. ",
        "conclusions": "The prompt emission spectrum of a GRB is generally fitted with Band model (Band et al. 1993). The empirical correlations, we have discussed, relies on the fact that the overall and instantaneous spectrum is Band like. Band model describes a non-thermal model. There are alternative models, e.g., black body with a powerlaw (Ryde 2004). Recently, we have analyzed the brightest GRBs, having separable pulses, namely GRB 081221 and GRB 090618, and found that the Band model is adequate in the falling part of a pulse. But, a different model is preferred in the rising part (Basak \\& Rao 2013). Hence, one should take into account these findings while describing the emission mechanism of GRB and thereby using the pulses for cosmological purpose."
    },
    "1402/1402.4119_arXiv.txt": {
        "abstract": "We report a weak line at $3.52\\pm 0.02$~keV in X-ray spectra of M31 galaxy and the Perseus galaxy cluster observed by MOS and PN cameras of XMM-Newton telescope.  This line is not known as an atomic line in the spectra of galaxies or clusters. It becomes stronger towards the centers of the objects; is stronger for Perseus than for M31; is absent in the spectrum of a deep ``blank sky'' dataset.  Although for each object it is hard to exclude that the feature is due to an instrumental effect or an atomic line, it is consistent with the behavior of a dark matter decay line.  Future (non-)detections of this line in multiple objects may help to reveal its nature. ",
        "introduction": "Significances quoted in the main body of the paper (Table~I) reflect the local significance of the signal. Since the position of the line is unknown \\textit{a priori} we need to take into account the probability of falsely detecting a statistical fluctuation of equal or higher significance at any position in the entire fitting range (2.0--8.0~keV). In addition, having found a signal in the same energy bin in three separate datasets, we compute this global significance taking into account the probability of such signals showing in the same resolution element by chance. Given the local significance of the signal in each dataset (based on the $\\Delta\\chi^2$ values and the number of degrees of freedom), and the number of independent resolution elements, we can determine the global significance of the combination of all signals. The number of independent resolution elements, $N_E$, for our datasets is about 40 (6~keV energy range divided by 150 eV --- average energy resolution of the \\xmm).  The global significance per dataset is computed from the two-sided p-value $p_i$ (directly related to the number of $\\sigma$ of the signal) by multiplying by $N_E$ (see Table~\\ref{tab:significance}). We took a ``two-sided'' p-value to take into account both positive and negative residuals.  The combined global significance then is \\begin{equation} \\label{A1} \\frac{\\prod_i p_i N_E}{N_E^{N_d-1}} = 1.1\\cdot 10^{-5} \\end{equation} where $N_d = 3$ is the number of datasets. This corresponds to a false detection probability for the combination dataset of \\textbf{0.0011\\%}. Converted to the significance this p-value gives \\textbf{$4.4\\sigma$ global significance}.  Alternatively, we could have taken into account only probability of positive fluctuations (so ``two-sided'' p-values in the Table~\\ref{tab:significance} should be divided by 2).  Using the same formula~\\eqref{A1} we would obtain $4.7\\sigma$ global significance.  Introducing systematic uncertainties into all our datasets at the level of $\\sim 1$\\%, the local significances drop by about $1\\sigma$ each.    \\begin{table*}[t!] \\centering \\begin{tabular}{l|c|c|c|c|c|c|c} \\hline Dataset & $\\Delta \\chi^2$ & d.o.f. & local significance & local p-value & false detection probability & global significance\\\\ \\hline M31-oncen (MOS)& 13 & 2 & 3.18$\\sigma$ & $1.5\\cdot 10^{-3}$ & 0.06 & 1.89$\\sigma$ \\\\ Perseus (MOS) & 9.1 & 2 & 2.56$\\sigma$ & $1.05\\cdot 10^{-2}$ & 0.42 & 0.81$\\sigma$ \\\\ Perseus (PN) & 8 & 2 & 2.36$\\sigma$ & $1.83 \\cdot 10^{-2}$ & 0.73 & 0.35$\\sigma$\\\\ \\hline All combined & & & & & $1.1 \\cdot 10^{-5}$ & 4.4$\\sigma$ \\\\ \\hline \\end{tabular} \\caption{\\label{tab:significance} Table of significances per dataset. Quoted p-values refer to the two-sided case (one-sided p-values are half of the two-sided ones). The false detection probability refers to the probability of falsely detecting a signal in that dataset like the one under consideration or stronger at any energy in the range considered. The global significance was converted from the false detection probability per dataset. The combined false detection probability and global significance of these three datasets is also given (computed from the individual detections, not from a single combined dataset).} \\end{table*}    \\begin{figure}[t!] \\centering \\includegraphics[width=.45\\textwidth]{mos-area.pdf} ~\\includegraphics[width=.45\\textwidth]{pn-area.pdf} \\caption{Exposure averaged effective area of the XMM MOS camera for the combination of observations of Perseus galaxy cluster, M31 and blank-sky (\\textit{left panel}). For Perseus galaxy cluster we also show the exposure averaged PN camera's effective area (\\textit{right panel}).} \\label{fig:aeff} \\end{figure}      \\setcounter{magicrownumbers}{0}  \\begin{table*}[!p] \\centering \\begin{tabular}[c]{r|c|c|c|c|c} \\hline &ObsID & Off-axis angle & Cleaned exposure & FoV  [arcmin$^2$] & F$_{in}$-F$_{out}$ \\\\ &&arcmin & MOS1/MOS2 [ksec] & MOS1/MOS2 & \\\\ \\hline \\rownumber & \\texttt{0305690301} & 22.80 & 18.6 / 18.6 & 473.6 / 574.3 & 1.266 / 1.340 \\\\ \\rownumber & \\texttt{0085590201} & 25.01 & 40.1 / 40.5 & 564.6 / 572.1 & 1.290 / 1.336 \\\\ \\rownumber & \\texttt{0204720101} & 27.87 & 14.1 / 14.5 & 567.7 / 574.5 & 2.373 / 2.219 \\\\ \\rownumber & \\texttt{0673020401} & 29.48 & 15.6 / 17.6 & 479.6 / 574.0 & 1.318 / 1.331 \\\\ \\rownumber & \\texttt{0405410201} & 29.52 & 16.1 / 16.6 & 480.8 / 573.9 & 1.354 / 1.366 \\\\ \\rownumber & \\texttt{0305690101} & 29.54 & 25.1 / 25.4 & 476.0 / 573.5 & 1.231 / 1.247 \\\\ \\rownumber & \\texttt{0405410101} & 31.17 & 15.8 / 16.8 & 481.8 / 572.9 & 1.235 / 1.195 \\\\ \\rownumber & \\texttt{0305720101} & 31.23 & 11.5 / 11.8 & 476.8 / 573.9 & 1.288 / 1.296 \\\\ \\rownumber & \\texttt{0673020301} & 36.54 & 13.9 / 15.4 & 485.4 / 573.8 & 1.211 / 1.304 \\\\ \\rownumber & \\texttt{0305690401} & 36.75 & 25.9 / 26.0 & 479.1 / 573.8 & 1.158 / 1.156 \\\\ \\hline \\rownumber & \\texttt{0305720301} & 41.92 & 16.7 / 17.5 & 464.7 / 573.6 & 1.433 / 1.447 \\\\ \\rownumber & \\texttt{0151560101} & 47.42 & 23.7 / 23.6 & 572.1 / 573.6 & 1.294 / 1.206 \\\\ \\rownumber & \\texttt{0673020201} & 53.31 & 22.8 / 23.4 & 479.5 / 573.9 & 1.262 / 1.228 \\\\ \\rownumber & \\texttt{0204720201} & 54.11 & 22.4 / 22.9 & 564.0 / 573.2 & 1.153 / 1.195 \\\\ \\hline \\rownumber & \\texttt{0554500801} & 95.45 & 15.0 / 15.3 & 480.8 / 572.7 & 1.098 / 1.113 \\\\ \\rownumber & \\texttt{0306680301} & 101.88 & 12.3 / 13.0 & 468.1 / 574.0 & 1.177 / 1.089 \\\\ \\hline \\end{tabular} \\caption{Parameters of the \\xmm spectra of the Perseus cluster used in our analysis.  The observations are sorted by the off-axis angle from the center of the Perseus cluster. Two  central observations (ObsIDs  \\texttt{0305780101} and  \\texttt{0085110101}) were not included in the analysis to avoid modeling of the emission from the core of the Perseus cluster. Notice that only  these two central observations were used in~[14], therefore our dataset and that of~[14] is independent from eachother. The difference in FoVs between MOS1 and MOS2 cameras is due to the loss CCD6 in MOS1 camera. The  parameter F$_{in}$-F$_{out}$ (last column) estimates the presence of residual soft protons according to the procedure of~[21]. Note, however, that for the bright extended sources, such an estimate is not appropriate, see~\\protect\\url{http://xmm2.esac.esa.int/external/xmm_sw_cal/background/epic_scripts.shtml} for details). Horizontal lines shows how we group observations for building the surface brightness profile of the line (as shown in Fig.~2, right panel in the main text).} \\label{tab:perseus-xmm-observations} \\end{table*}  \\begin{table}[p!] \\centering \\begin{tabular}{c|c|c} \\hline Range of offsets & Exposure [ksec] & Flux [$\\unit{cts/sec/cm^2}$] \\\\ \\hline $23$ -- $37'$ & 400 & $13.8 \\pm 3.3$\\\\ $42'$ -- $54'$ & 230 & $8.3 \\pm 3.4$\\\\ $96'$ -- $102'$& 56 & $4.6 \\pm 4.6$\\\\ \\hline \\end{tabular} \\caption{Definitions of the radial bins used for the data analysis of the Perseus cluster.} \\label{tab:radial} \\end{table}   \\begin{figure*}[t!] \\centering \\includegraphics[width=0.45\\textwidth]{perseus_offcen_mos12-all.pdf}~% \\includegraphics[width=0.45\\textwidth]{perseus_offcen_mos12-zoom3-4.pdf} \\caption{\\textit{Left:} Folded count rate (top) and residuals (bottom) for the combined spectrum of 16 observations of MOS cameras (listed in the Tabel~\\protect\\ref{tab:perseus-xmm-observations}) of the Perseus cluster. Statistical Y-errorbars on the top plot are smaller than the point size. The line around 3.5~keV is \\emph{not added}, hence the group of positive residuals. \\textit{Right}: zoom onto the line region. The spectrum is shown in the detector restframe, therefore the line is shifted left according to the Perseus redshift.} \\label{fig:perseus} \\end{figure*} ",
        "conclusions": ""
    },
    "1402/1402.1350_arXiv.txt": {
        "abstract": "{We provide flux-limited and volume-limited galaxy group and cluster catalogues, based on the spectroscopic sample of the SDSS data release 10 galaxies. We used a modified friends-of-friends (FoF) method with a variable linking length in the transverse and radial directions to identify as many realistic groups as possible. The flux-limited catalogue incorporates galaxies down to $m_r=17.77~\\mathrm{mag}$. It includes 588193 galaxies and 82458 groups. The volume-limited catalogues are complete for absolute magnitudes down to $M_{r,\\mathrm{lim}} = -18.0$, $-18.5$, $-19.0$, $-19.5$, $-20.0$, $-20.5$, and $-21.0$; the completeness is achieved within different spatial volumes, respectively. Our analysis shows that flux-limited and volume-limited group samples are well compatible to each other, especially for the larger groups/clusters. Dynamical mass estimates, based on radial velocity dispersions and group extent in the sky, are added to the extracted groups. The catalogues can be accessed via \\texttt{http://cosmodb.to.ee} and the Strasbourg Astronomical Data Centre (CDS).} ",
        "introduction": "Galaxies, as well as stars or people, tend to gather in pairs, groups of several members, or even larger conglomerates. It is only natural that an individual surrounded by companions follows a somewhat different evolutionary path and develops a different appearance and inner qualities than a loner; the dependence of galaxy properties on the environment was shown already decades ago \\citep{Einasto:74,Oemler:74,Davis:76,Dressler:80,Postman:84}. Therefore, catalogues of galaxy groups and clusters provide an unlimited data source for a wide range of astrophysical and cosmological applications, as illustrated below with just a random pick of the latest studies.  Among the most popular applications, galaxy group and cluster catalogues are used for quantifying the neighbourhood of galaxies in studies of the environmental dependencies of galaxy properties. The published catalogues have enabled a multitude of studies about the dependencies of the morphology, structure, gas content, star formation rate, merger rate, etc. of galaxies on their local environment \\citep[some most recent examples include][]{Tempel:09, Carollo:13,Catinella:13,Hess:13,Kaviraj:13,Lackner:13,Peng:13}.  Besides, group and cluster catalogues can be used to distinguish central galaxies from their satellites \\citep{Lacerna:13,Wetzel:13,Yang:13,Li:14}. They also simplify the selection and studying of the properties of specific galaxy systems, be it galaxy pairs, compact groups, loose groups, etc. Assuming that galaxy groups and clusters inhabit common dark matter haloes, catalogues can be used to probe dark matter haloes of different mass and to seek correlations between the mass and the galaxy content of the haloes \\citep{Carollo:13,Hou:13,Huertas:13,McGee:13,Wetzel:13,Yang:13}, or to study the general group properties as a function of group mass \\citep{Hearin:13, Budzynski:14} or as a function of the large-scale environment \\citep{Luparello:13,Einasto:14}. By comparing the properties of observed groups and clusters to those produced in simulations, our understanding of structure formation can be validated \\citep{Nurmi:13}.  Yet another broad field of applications for group and cluster catalogues is their usage for mapping the underlying cosmic web, e.g. by extracting large-scale filaments \\citep{Zhang:13,Alpaslan:13}, or the largest known density enhancements, galaxy superclusters \\citep{Zucca:93, Einasto:94, Einasto:01} from the spatial distribution of groups and clusters. Further, the properties of groups in superclusters, and galaxies in them can be studied \\citep{Einasto:03a,Einasto:08, Einasto:11, Einasto:12, Lietzen:12, Krause:13}.  From the examples above, it is evident that catalogues of galaxy groups and clusters are in demand. But how to define a galaxy group and how to delineate it from galaxy redshift survey data? No straightforward recipe exists, different approaches would be valid for different science goals. In simulations we can see that gravitationally bound galaxy systems are linked together by an underlying dark matter halo, thus a good approach for defining a galaxy group or cluster would be through the existence of a common dark matter halo. However, for observational datasets, this method is of very little practical value. Instead, the friends-of-friends (FoF) algorithm has remained the most frequently applied means of identifying groups and clusters in galaxy redshift data ever since its introduction \\citep{Turner:76,Press:82}. Today, many extensive FoF group and cluster catalogues are available, varying in the sample data and details of the group finder algorithm \\citep[e.g. ][]{Tucker:00,Eke:04,Tago:06,Tago:08}.  The FoF method uses galaxy distances as the main basis of grouping, and is thus relatively simple and straightforward, while the membership of the produced groups is rather uncertain. Thus several other techniques for group and cluster extraction have been developed, well reviewed by \\cite{Gal:06} and recently applied by \\citet{Yang:05}, \\citet{Koester:07}, \\citet{Yang:07}, \\citet{Hao:10}, \\citet{Makarov:05}, \\citet{Wen:12}, and \\citet{MunosCuartas:12}. However, whether any method is more reliable than the others is still largely a matter of taste and debate \\citep{Old:14}.  In \\cite{Tempel:12a}, we constructed a flux-limited FoF group and cluster catalogue for galaxies in the Sloan Digital Sky Survey \\citep[SDSS,][]{York:00} data release 8 \\citep[DR8;][]{Aihara:11}. Here we present an update of this catalogue, based on SDSS~DR10 \\citep{Ahn:13}. In addition to a flux-limited catalogue, we present volume-limited catalogues, valid within different galactic absolute luminosity bins (and thus being complete within different spatial volumes). We provide dynamical mass estimates for the detected galaxy systems, using the measured radial velocities and group extent in the sky. The catalogues also contain an rough estimate of the expected total luminosity of each group, assuming that some of the group members are not included on the sample due to observational limitations.  Throughout this paper we assume the Wilkinson Microwave Anisotropy Probe (WMAP) cosmology: the Hubble constant $H_0 = 100\\,h\\ \\mathrm{km\\,s^{-1}Mpc^{-1}}$, the matter density $\\Omega_\\mathrm{m}=0.27$ and the dark energy density $\\Omega_\\Lambda=0.73$ \\citep{Komatsu:11}. ",
        "conclusions": "We have updated and improved our previous galaxy group catalogues, constructed on the basis of the FoF method, following a similar procedure as used in the previous papers \\citep{Tago:08,Tago:10,Tempel:12a}. The group finding method is applied to flux- and volume-limited samples drawn from the SDSS main contiguous area, covering 7221 square degrees in the sky. In addition to the SDSS spectroscopic redshifts, the galaxy sample is complemented with redshifts from the 2MRS, 2dFGRS and RC3 catalogues.  As an important addition, we have estimated the group/cluster dynamical masses using the virial theorem. We have shown that the groups extracted from the flux- and volume-limited samples are well compatible with each other. We have calculated the mass functions for the volume-limited samples. At the massive end, the functions are in agreement with previous observational estimates \\citep{Rines:07}.  As a next step, we plan to use the new catalogue to map the large-scale galactic filaments \\citep[similarly to][]{Tempel:14} and to construct a catalogue of superclusters \\citep[similarly to][]{Liivamagi:12}.  \\begin{figure} \\centering \\includegraphics[width=88mm]{fig13.pdf} \\caption{Group mass function for various volume-limited samples (solid lines) together with 95\\% confidence limits (shaded regions). The mass function is calculated using groups with three or more members. Black and red points and error bars are mass function from \\citet{Rines:07} using virial masses and caustic masses, respectively (the error bars show 68\\% uncertainties).} \\label{fig:gr_massfun} \\end{figure}"
    },
    "1402/1402.4743_arXiv.txt": {
        "abstract": "Modeling the spectral emission of low-charge iron group ions enables the diagnostic determination of the local physical conditions of many cool plasma environments such as those found in H {\\sc ii} regions, planetary nebulae, active galactic nuclei etc. Electron-impact excitation drives the population of the emitting levels and, hence, their emissivities. By carrying-out Breit-Pauli and intermediate coupling frame transformation (ICFT) $R$-matrix calculations for the electron-impact excitation of Fe$^{2+}$ which both use the exact same atomic structure and the same close-coupling expansion, we demonstrate the validity of the application of the powerful ICFT method to low-charge iron group ions. This is in contradiction to the finding of Bautista et al. [\\apjl, 718, L189, (2010)]  who carried-out ICFT and Dirac $R$-matrix calculations for the same ion. We discuss possible reasons. ",
        "introduction": "\\label{sec:intro}    \\begin{table*} \\caption{Lowest 17 energy levels (Ryd) of the $3d^6$ ground configuration of Fe$^{2+}$. \\label{tab:energies}} \\begin{small} \\begin{center} \\begin{tabular}{lllll} \\tableline \\tableline Term/Level & Observed\\tablenotemark{a} & AS\\tablenotemark{b} & GRASP0\\tablenotemark{b}  & GRASP0\\tablenotemark{c}  \\\\ \\tableline $^5D_4$ &   0.0        & 0.0      & 0.0       & 0.0      \\\\ $^5D_3$ &   0.003975   & 0.004001 & 0.003408  & 0.003613 \\\\ $^5D_2$ &   0.006733   & 0.006803 & 0.005823  & 0.006162 \\\\ $^5D_1$ &   0.008497   & 0.008602 & 0.007381  & 0.007803 \\\\ $^5D_0$ &   0.009361   & 0.009483 & 0.008147  & 0.008608 \\\\ $^3P2_2$&   0.176830   & 0.18434  & 0.19477   & 0.19564  \\\\ $^3P2_1$&   0.188527   & 0.19636  & 0.20543   & 0.20690  \\\\ $^3P2_0$&   0.193266   & 0.20122  & 0.20994   & 0.21161  \\\\ $^3H_6$ &   0.182719   & 0.20976  & 0.22119   & 0.22244  \\\\ $^3H_5$ &   0.184995   & 0.21173  & 0.22282   & 0.22416  \\\\ $^3H_4$ &   0.186645   & 0.21370  & 0.22444   & 0.22588  \\\\ $^3F2_4$&   0.195578   & 0.21026  & 0.22199   & 0.22313  \\\\ $^3F2_3$&   0.197744   & 0.21313  & 0.22439   & 0.22571  \\\\ $^3F2_2$&   0.199177   & 0.21468  & 0.22572   & 0.22712  \\\\ $^3G_5$ &   0.223796   & 0.24873  & 0.26285   & 0.26437  \\\\ $^3G_4$ &   0.227278   & 0.25202  & 0.26558   & 0.26729  \\\\ $^3G_3$ &   0.229114   & 0.25367  & 0.26700   & 0.26878  \\\\ \\tableline \\tablenotetext{1}{\\cite{NIST}} \\tablenotetext{2}{Present work.} \\tablenotetext{3}{From \\cite{Bau10}}  \\end{tabular} \\end{center} \\end{small} \\end{table*}   The electron-impact excitation of low-charge ion group ions is an important process in many cool astrophysical plasmas. The subsequent spectral emission can be used to diagnose the physical conditions of H {\\sc ii} regions, planetary nebulae, active galactic nuclei and many stars --- see for example \\cite{Joh00,Ves01,Est04,Mes09,Meh10,Mad12,Zha12,Mcc13}.  In order to have confidence in the temperatures, densities, and chemical compositions deduced from modeling such emission it is important that the uncertainties in the atomic data be well understood and their effects modeled. Energy levels can be compared with observed. Radiative data can usually be computed with increasingly large configuration interaction expansions to demonstrate convergence to a sufficient degree. The case of electron-impact excitation is more problematic. The complexity of the open $3d$-subshell leads to a large number of low-lying target levels which should be coupled together in a scattering calculation \\citep{Burk94}. The $R$-matrix method is paramount here. It has been advantageous to introduce some approximations into the scattering problem which leads to the highly efficient and widely used intermediate coupling frame transformation (ICFT) $R$-matrix method \\citep{Griff98}.  \\cite{Bau10} recently modeled Fe {\\sc iii} line intensities from the Orion nebulae which had been observed by \\cite{Mes09}. They found much better accord between theory and observation when they used excitation data which they calculated with the Dirac Atomic $R$-matrix code (DARC) methodology \\citep{Ait96, Nor04} than that they calculated with the ICFT method. Their ICFT Maxwellian-averaged collision strengths for transitions from the ground level to all other levels of the $3d^6$ configuration differed by up to a factor 3 (both larger and smaller) from their DARC ones at $10^{4}$~K. They attributed this difference ``to the limited treatment of spin-orbit effects in near-threshold resonances'' by the ICFT method. This surprising finding has implications not just for the validity and use of the ICFT method but also the $R$-matrix II \\citep{Burk94} FINE method \\citep{Burk10} with which it shares some similarities.  There are a number of issues which can be raised about the comparison made by \\cite{Bau10}.  It has long been known that to be able to deduce anything meaningful about two scattering methods from a comparison of collision data it is necessary that the differences in the atomic structure be kept to an absolute minimum. \\cite{Bau10}  used a much larger configuration interaction (36 configurations) for their ICFT atomic structure compared to that for their DARC one (8 configurations). They also used a number of pseudo orbitals in their ICFT target description while the DARC ones were purely spectroscopic --- pseudo-orbitals lead to pseudo-resonances. They retained 283 levels in their ICFT close-coupling expansion but 322 levels in their DARC one.  We report-on the results of parallel ICFT and BPRM $R$-matrix calculations for Fe$^{2+}$ that use the exact same atomic structure (an 8 configuration interaction expansion) and the same close-coupling expansion (322 levels). We find excellent agreement between the two ($<5$\\% where \\cite{Bau10}  found factors of 3). We also find excellent agreement with a similar DARC calculation which we have carried out. Many of the differences found by \\cite{Bau10}  disappear once we correctly label their DARC transitions. Some differences remain with their ICFT results and for which the source is likely one of the possible causes listed in the preceding paragraph.  The rest of this paper is structured as follows: in Section 2 we describe the key differences between the different $R$-matrix methods we use; in Section 3 we give details of both the atomic structure and atomic collision calculations; in Section 4 we present our results and discuss their agreement with and differences from the results of \\cite{Bau10}; and in Section 5 we draw some key conclusions. ",
        "conclusions": "We have carried-out ICFT and Breit-Pauli\\footnote{A full Breit-Pauli dataset for all 322 levels (energies, radiative transition probabilities and effective collision strengths) is available in electronic form from the APAP website (\\url{www.apap-network.org}).} $R$-matrix calculations for the electron-impact excitation of Fe$^{2+}$ which use the exact same atomic structure and the same close-coupling level expansion.  The results demonstrate that the ICFT $R$-matrix method can be expected to provide accurate effective collision strengths for near neutral iron group ions which are well within the uncertainties which exist due to the accuracy of the representation of their atomic structure.  The advantage of the ICFT approach over the Breit-Pauli and Dirac $R$-matrix approaches is that calculations which require the resources of massively parallel supercomputer centers can typically be carried-out on small local clusters. This facilitates the study of near neutral iron group ions which are omnipresent in cool astrophysical plasmas and which shine a light on their physical conditions."
    },
    "1402/1402.1166_arXiv.txt": {
        "abstract": "We use the SDSS Stripe 82 to empirically quantify the stellar-mass and black-hole growth triggered by minor mergers in local spiral (disk) galaxies. Since major mergers destroy disks and create spheroids, morphologically disturbed spirals are likely remnants of minor mergers. Disturbed spirals exhibit enhanced specific star formation rates (SSFRs), the enhancement increasing in galaxies of `later' morphological type (which have more gas and smaller bulges). By combining the SSFR enhancements with the fraction of time spirals spend in this `enhanced' mode, we estimate that $\\sim$40\\% of the star formation in local spirals is directly triggered by minor mergers. The disturbed spirals also exhibit higher nuclear-accretion rates, implying that minor mergers enhance the growth rate of the central black hole. However, the specific accretion rate shows a lower enhancement than that in the SSFR, suggesting that the coupling between stellar-mass and black-hole growth is weak in minor-merger-driven episodes. Given the significant fraction of star formation that is triggered by minor mergers, this weaker coupling may contribute to the large intrinsic scatter observed in the stellar vs. black-hole mass relation in spirals. Combining our results with the star formation in early-type galaxies -- which is minor-merger-driven and accounts for $\\sim$14\\% of the star formation budget -- suggests that \\emph{around half of the star formation activity in the local Universe is triggered by the minor-merger process.} ",
        "introduction": "The standard $\\Lambda$CDM model of galaxy formation postulates a hierarchical growth of structure over cosmic time \\citep[e.g.][]{White1978,Somerville1999,Cole2000,Hatton2003}. In this paradigm, galaxies and their host dark-matter halos grow, both through mergers with systems of similar mass (`major mergers'), and via the accretion of smaller objects (`minor mergers'). A natural consequence of the shape of the observed luminosity function -- in which smaller galaxies far outnumber their more massive counterparts \\citep[e.g.][]{Cole2001,Blanton2001} -- is that minor mergers are significantly more common than major interactions.  While past studies of merging have largely explored the effects of major mergers \\citep[e.g.][]{Darg2010a,Darg2010b,Toomre1972,Hernquist1989,Barnes1992a,Jesseit2007}, a growing literature is highlighting the important role of minor mergers (mass ratios $\\lesssim$ 1:4) in influencing massive-galaxy evolution. Minor merging is a frequent process, that is both predicted \\citep[e.g.][]{Maller2006,Stewart2008,Fakhouri2008,Kaviraj2009} and observed \\citep[e.g.][]{Lin2004,Jogee2009,Lopez2010} to be at least $\\sim$3-4 times more common than major interactions at late epochs. While each individual minor interaction may have a relatively small effect on the massive galaxy, it is becoming increasingly clear that the \\emph{cumulative} impact of minor mergers over cosmic time is significant, both in terms of the stellar content of massive galaxies and their structural properties.  \\begin{figure*} \\begin{minipage}{172mm} \\begin{center} \\includegraphics[width=\\textwidth]{minor_merger_remnants4} \\caption{Examples of standard-depth, multi-colour images (left) and their deeper $r$-band Stripe 82 counterparts (right) of spiral minor-merger remnants. We show examples of early-type (row 1, left), Sa (row 1, middle), Sb (row 1, right), Sc (row 2, left) and Sd (row 2, middle) galaxies. The morphological disturbances are typically clearer in the Stripe 82 imaging, making these deeper images necessary for this exercise. While ETGs are not considered in this study, we show an example of a disturbed ETG for comparison to the spirals. Note that the images may look better on screen than in print.} \\label{fig:hubble_sequence} \\end{center} \\end{minipage} \\end{figure*}  The role of this process in fuelling star formation in massive galaxies has recently been explored via ultraviolet (UV) studies of early-type galaxies (ETGs). Survey-scale UV data, from GALEX at $z\\sim0$ (e.g. Yi et al. 2005; Kaviraj et al. 2007) and deep optical surveys at intermediate redshift \\cite[e.g.][]{Kaviraj2008b}, have revealed widespread star formation in ETGs at late epochs. A strong correlation is observed between the presence of morphological disturbances and blue UV colours in nearby ETGs (Kaviraj et al. 2011; see also Schweizer et al. 1990, 1992), indicating that the star formation is driven by mergers, and not by processes like stellar mass loss or accretion that would leave the stellar morphology of the system undisturbed (Kaviraj et al. 2011, see also Crockett et al. 2011). However, the major-merger rate at late epochs ($z<1$) cannot satisfy the number of disturbed ETGs, suggesting that \\emph{minor} mergers dominate the star formation in these objects (Kaviraj et al. 2011, see also Shabala et al. 2012). Note that these studies primarily sample low-density environments, making these results relevant largely to ETGs that inhabit groups and the field (and not dense clusters). However, since clusters are typically hostile to star formation \\citep[e.g.][]{Dressler1984,Moore1999,Kimm2011}, the bulk of the ETG star formation budget is expected to reside in low-density environments.  While past empirical work on minor-merger-driven star formation has focused on ETGs, this process will clearly affect massive galaxies regardless of morphology. \\citet[][K14 hereafter]{Kaviraj2014} has leveraged the ETG studies mentioned above to derive a lower limit for the fraction of star formation in spiral galaxies that is also likely triggered by minor mergers. The K14 results indicate that \\emph{at least} a quarter of the star formation budget in local spirals is attributable to this process, yielding a lower limit for the minor-merger-driven fraction of \\emph{cosmic} star formation of $\\sim$35\\%. The observed correlation between black hole and galaxy stellar mass \\citep[e.g.][]{Gultekin2009} further implies that a similar fraction of black-hole growth may also be induced by this process.  The work on minor-merger-driven star formation is mirrored by studies of the size evolution of massive galaxies over cosmic time. Recent empirical work has demonstrated that the effective radii of massive galaxies increase, on average, by factors of 3-5 over the lifetime of the Universe \\citep[e.g.][]{Daddi2005,Trujillo2006,VD2008,Buitrago2008,Saracco2009,Cimatti2012,Newman2012,Ryan2012,Huertas-Company2012}. While secular mechanisms, such as adiabatic expansion driven by stellar mass loss or strong AGN feedback have been proposed to explain this evolution \\citep[e.g.][but see Trujillo et al. 2011 and Bluck et al. 2012]{Fan2008,Damjanov2009}, consensus favours size evolution via mergers, with the bulk of the increase attributed to minor mergers over cosmic time \\citep[e.g.][]{Khochfar2006,Bournaud2007,Naab2009,Nipoti2009,Oser2012}.  While it is increasingly well-established that the minor-merger process can significantly influence the star-formation and structural properties of massive galaxies, its overall impact on cosmic star formation and black hole growth remains poorly understood. Since spirals dominate the star-formation budget at late epochs (K14), this translates into a need to understand, in detail, the impact of minor mergers on the spiral population across cosmic time. While K14 has derived a lower limit ($\\sim$25\\%) for the fraction of star formation in spirals that is triggered by this process at the present day, an estimate of its global impact on galaxy evolution demands a more accurate measurement of this fraction, beyond the lower limit derived by K14.  A burgeoning literature has now begun probing the impact of minor mergers on spiral galaxies. Empirical studies of spiral minor-merger \\emph{remnants} have largely focused on individual case studies of very nearby galaxies. These efforts typically demonstrate enhancements in star formation and, in some cases, nuclear activity \\citep[see e.g.][]{Smith1996,Knapen2004,Mazzuca2006}, indicating that the minor-merger process is indeed influencing the growth of both the galaxy and its central black hole. {\\color{black}Studies of minor galaxy \\emph{pairs} -- i.e. systems where the two objects have not coalesced -- also show enhancements, especially when the system is close to coalescence (e.g. separations $<10$ kpc), although the effect is much weaker than in major mergers \\citep{Li2008,Ellison2008,Robaina2009,Scudder2012,Patton2013}.} The triggering of star formation in the massive progenitor of a minor merger is expected to peak at the `remnant' stage, during the coalescence of the satellite, when the tidal forces on the disk are highest \\citep[see e.g.][]{Mihos1994,Kaviraj2009} - statistical studies of minor-merger \\emph{remnants} are, therefore, very desirable for gaining insights into this process.  \\begin{figure*} \\begin{minipage}{172mm} \\begin{center} \\includegraphics[width=\\textwidth]{relaxed_disks2} \\caption{Examples of standard-depth, multi-colour (left) images and their deeper $r$-band Stripe 82 counterparts (right) of relaxed spiral galaxies. Note the lack of tidal disturbances compared to the objects shown in Figure 1 above. We show examples of early-type (row 1, left), Sa (row 1, middle), Sb (row 1, right), Sc (row 2, left) and Sd (row 2, middle) galaxies. While ETGs are not considered in this study, we show an example of an ETG for comparison to the spirals. Note that the images may look better on screen than in print.} \\label{fig:minor_merger_remnants} \\end{center} \\end{minipage} \\end{figure*}  In a similar vein, while past theoretical efforts have largely explored the impact of \\emph{major} mergers, recent work has begun to study the role of minor interactions in spiral galaxies. In qualitative agreement with the empirical case studies mentioned above, theoretical efforts do find star-formation enhancement in minor-merger remnants. The enhancement typically anti-correlates with the system mass ratio \\citep[e.g.][]{Hernquist1989,Mihos1994,Hernquist1995,Cox2008}, but the actual increase in star-formation efficiency depends critically on the specific theoretical prescriptions utilized in individual simulations.  Since the enhancement of star formation occurs at the `remnant' stage of the process i.e. when the smaller companion is coalescing with the primary, an \\emph{empirical} survey-scale study of spiral minor-merger remnants becomes a compelling exercise. In this paper we present such a study, whose principal aims are to (1) quantify the enhancement of star formation due to minor mergers in spiral galaxies (2) estimate the fraction of total star formation in spirals, and in massive galaxies in general, that is triggered by minor mergers and (3) quantify the enhancement of black hole growth in spirals due to the minor-merger process. This study significantly extends the K14 analysis (which derived a robust lower limit for point [2] above), with the overall objective of gauging the global impact of minor mergers on the evolution of the spiral galaxy population.  This paper is organized as follows. In Section 2, we describe the galaxy sample that underpins this study and the selection of minor-merger remnants via visual inspection. In Sections 3 and 4, we explore the enhancement of star formation in spirals due to minor mergers and use these results to estimate the fractions of the spiral and cosmic star-formation budgets that are triggered by the minor-merger process. In Section 5, we explore minor-merger-driven enhancement of nuclear accretion in spiral galaxies. We summarize our findings in Section 6. Throughout this study we use the \\emph{WMAP3} cosmological parameters \\citep{Komatsu2011}: $\\Omega_m = 0.241$, $\\Omega_{\\Lambda} = 0.759$, $h=0.732$, $\\sigma_8 = 0.761$. ",
        "conclusions": "We have used a sample of bright ($r<16.8$), local ($z<0.07$) galaxies from the SDSS Stripe 82 to probe the role of minor mergers in driving stellar mass and black hole growth in nearby massive spiral (disk) galaxies. Our study has been based on a sample of galaxies from K14, that have been visually classified into standard morphological classes (E/S0, Sa, Sb, Sc, Sd, etc.) using both colour images from the SDSS DR7 and their deeper $r$-band counterparts from the Stripe 82.  Since `major' (i.e. equal or nearly equal-mass) mergers produce spheroids, an effective route to studying spiral galaxies that are minor-merger remnants is to use spirals that are morphologically disturbed. The disturbed morphology indicates a recent interaction, and the continued presence of a disk indicates that the interaction was not a major merger (since this would have destroyed the disk and created a spheroid). Using the DR7 and Stripe 82 images, we have flagged spiral galaxies in the K14 sample that are morphologically disturbed, thus selecting the nearby spiral minor-merger remnant population in the SDSS Stripe 82 field.  We have used this sample to quantify the stellar mass growth in spiral galaxies that is plausibly triggered by minor mergers. As indicated by Eqn. 2 and described in Section 3.1, the proportion of minor-merger-driven star formation depends both on the enhancement of star formation during the remnant phase of a minor merger ($\\eta$) and the fraction of time galaxies spend in this enhanced star-formation mode (the `duty cycle', $D$). While we cannot measure these quantities for individual galaxies, we can calculate mean statistical estimates for $\\eta$ and $D$, if large populations of minor-merger remnants are available, as is the case in this study.  Assuming that the detectability of morphological disturbances does not evolve in the short redshift range studied here, $D$ can be estimated by the fraction of galaxies that are morphologically disturbed. $\\eta$ can be calculated using the ratio of the measured SSFRs in the relaxed spirals to that in the disturbed spirals.  The duty cycle is relatively insensitive to morphological class and ranges between 11 and 17\\%. The SSFRs are enhanced in the disturbed spirals, regardless of morphology, with the enhancement increasing in galaxies that have `later' morphological type. While disturbed Sa galaxies exhibit SSFR enhancements of a factor of $\\sim$2, this increases to a factor of $\\sim$6 for Sc and later morphological types. Our results indicate that the star-formation enhancements are higher in galaxies that host larger internal gas reservoirs and smaller bulges (that are less able to stabilize the disk against radial gas inflows), as might be expected from the recent literature.  The duty cycles and SSFR enhancements imply that the fraction of star formation driven by minor mergers range from $\\sim$27\\% in Sa galaxies to $\\sim$43-48\\% in Sb and later morphological types. Combining this with the proportion of the star formation budget hosted by each morphological class then yields \\emph{a total minor-merger-driven fraction of star formation in spirals of $\\sim$40\\%}. This is consistent with the results of K14, who estimated an empirical lower limit for this value of $\\sim$25\\%. Thus, while most of the star formation in today's spiral galaxies is unrelated to mergers, a significant fraction is attributable to the minor-merger process. It is worth noting that our analysis has been restricted to galaxies that show morphological disturbances. Thus, we have disregarded systems that might have undergone a recent minor merger, but where the mass ratio is not high to induce morphological disturbances at the depth of the Stripe 82 imaging. {\\color{black}We have also disregarded the contribution of the phase when minor galaxy pairs are close to coalescence but still on approach (although we have shown that this contribution is relatively insignificant)}. Nevertheless, both these points imply that the estimated value of 40\\% above may also be a lower limit. Combining our results with the star formation in early-type galaxies, which is dominated by minor mergers (at least in low-density environments) and accounts for 14\\% of cosmic star formation budget (K14), indicates that \\emph{around half the star formation activity at the present day is likely triggered by the minor-merger process.}  The observed correlation between galaxy stellar mass and the mass of the central black hole suggests that the enhanced star formation episodes triggered by minor mergers may also be accompanied by enhanced nuclear accretion. We have studied the specific accretion rate in `early' (Sa) and `late' (Sb and later types) spirals, by using the quantity L[OIII]/$\\sigma^4$, restricted to galaxies that are classified as Seyfert (in which the ionization is dominated by the central AGN). Mirroring the star formation results, the specific accretion rate is found to be enhanced in the disturbed spirals, with the enhancement higher in `later' morphological types. However, the enhancement in nuclear accretion is not as pronounced as that found in the star formation activity, suggesting that star formation and nuclear accretion may not be tightly coupled in these minor-merger-enhanced episodes. This relatively weak correlation may be partly responsible for the larger intrinsic scatter in the stellar vs. black-hole mass correlation in spiral galaxies compared to their early-type counterparts.  Combined with the recent study of K14, the results of this paper further highlight the important (and still poorly quantified) role of minor mergers in influencing stellar mass and black hole growth in massive galaxies. While a significant fraction of star formation in \\emph{nearby} spirals is triggered by minor mergers, the role of this process is likely to become more significant with increasing redshift, both due to the gradually increasing merger rate \\citep[e.g.][]{Lotz2011} and the higher gas fractions in massive galaxies at earlier epochs \\cite[e.g.][]{Tacconi2010}. At $z>1$, and certainly around the epoch of peak cosmic star formation ($z\\sim2$), a significant number of disk galaxies are clumpy, gravitationally unstable systems \\citep[see e.g.][]{Keres2009,Dekel2009b,Devriendt2010,Ceverino2010}, in which star formation could be efficiently triggered by minor mergers \\citep[e.g.][]{Birnboim2003,Keres2005,Dekel2009b}. Since disks form before bulges \\citep[e.g.][]{Hatton2003}, the lack of the stabilizing influence of a bulge may make these early disks even more susceptible to minor-merger-triggered star formation \\citep[e.g.][]{Mihos1994,Hernquist1995} compared to their modern-day counterparts.  Given the significance of this process in driving star formation, black-hole growth and size evolution in massive galaxies, studying the role of minor mergers across a large range in redshift is essential. The most detailed insights are likely to come from high-resolution spatially-resolved studies of local minor-merger remnants using the HST, that employ sensitive tracers of star formation like the UV. As noted in the introduction above, local star-forming ETGs are excellent initial laboratories for such spatially-resolved analyses that probe the star-formation laws associated with satellite accretion (Kaviraj et al., in prep). At intermediate redshift ($z<2$), the population of minor-merger remnants can be studied using the deep, high-resolution imaging offered by the HST \\citep[e.g.][]{Kaviraj2011}, in legacy fields such as CANDELS \\citep[e.g.][]{Grogin2011,Koekemoer2011}. At earlier epochs this exercise will likely require deep, high-resolution imaging from future instruments like LSST and the Extremely Large Telescopes. In any case, without a robust determination of minor-merger-driven stellar mass and black hole growth over cosmic time, our understanding of galaxy evolution is likely to remain incomplete."
    },
    "1402/1402.3952_arXiv.txt": {
        "abstract": "{ We review the quantum field theory description of Hawking radiation from evaporating black holes and summarize what is known about Hawking radiation from black holes in more than four space-time dimensions. In the context of the Large Extra Dimensions scenario, we present the theoretical formalism for all types of emitted fields and a selection of results on the radiation spectra. A detailed analysis of the Hawking fluxes in this case is essential for modelling the evaporation of higher-dimensional black holes at the LHC, whose creation is predicted by low-energy models of quantum gravity. We discuss the status of the quest for black-hole solutions in the context of the Randall-Sundrum brane-world model and, in the absence of an exact metric, we review what is known about Hawking radiation from such black holes. } ",
        "introduction": "\\label{sec:intro}  Hawking radiation \\cite{Hawking:1975sw} is one of the most important effects arising from quantum field theory in curved space, a semi-classical approach to quantum gravity. In this framework space-time is described by a classical geometry, governed by the Einstein equations (or an alternative classical theory of gravity). The behaviour and propagation of quantum fields on a fixed (but not necessarily stationary) space-time is then studied. Hawking radiation is thermal in nature, giving non-extremal black holes an intrinsic temperature proportional to the surface gravity of the event horizon. For a Schwarzschild black hole in asymptotically flat space, the specific heat is negative, so that the temperature increases as the black hole evaporates, leading to a black hole explosion. The ultimate fate of the black hole depends on unknown details of quantum gravity, but black hole evaporation raises many deep questions about the nature of quantum gravity and the fundamental laws of physics (such as the information loss paradox, see, for example, \\cite{Hossenfelder:2009xq,Mathur:2009hf}). Here we will not address these important issues, but instead focus on the detailed properties of the Hawking radiation itself.  Hawking radiation from four-dimensional black holes in asymptotically flat space was studied in detail by Page \\cite{Page:1976df,Page:1976ki,Page:1977um}. Over the past fifteen or so years,  there has been great interest in higher-dimensional black holes. Within the context of classical general relativity, a menagerie of black-hole-like solutions of the Einstein equations has been discovered (see for example \\cite{Emparan:2008eg,Tomizawa:2011mc} for reviews). It is then natural to study the properties of Hawking radiation from higher-dimensional black holes.  This avenue of research gained much impetus from the exciting possibility of producing microscopic higher-dimensional black holes in high-energy collisions either at the LHC or in cosmic rays \\cite{Banks:1999gd,Dimopoulos:2001hw,Giddings:2001bu, Kanti:2004nr,Landsberg:2003br,Majumdar:2005ba,Park:2012fe,Webber:2005qa}. This is a prediction of higher-dimensional brane-world models \\cite{ArkaniHamed:1998rs,Antoniadis:1998ig,ArkaniHamed:1998nn,Randall:1999ee,Randall:1999vf} in which the energy scale of quantum gravity is much lower than the traditional value of $10^{19}$ GeV, and may be as low as the TeV-scale. If such a microscopic black hole is produced, it will initially be rapidly rotating and rather asymmetric. Its subsequent evolution can be modelled as four stages \\cite{Giddings:2001bu}: \\begin{itemize} \\item During the {\\em {balding phase}} the black hole sheds its asymmetries through the emission of gravitational radiation and also loses any gauge field charges arising from the particles which formed it. At the end of this stage the black hole is axisymmetric and still rapidly rotating. \\item The black hole then emits Hawking radiation, and loses both mass and angular momentum.  At the end of this {\\em {spin-down phase}} the black hole is no longer rotating. \\item Now with zero angular momentum, the black hole continues to radiate during the {\\em {Schwarzschild phase}}, shrinking as it loses mass. \\item During the final {\\em {Planck phase}} the semi-classical approximation for the Hawking radiation is no longer valid and the black hole emission depends on the details of quantum gravity. \\end{itemize} It is expected that the spin-down and Schwarzschild phases will dominate the life-time of the black hole. A detailed understanding of the Hawking radiation from higher-dimensional black holes in brane-world models is therefore necessary for simulating microscopic black hole events \\cite{Dai:2007ki,Frost:2009cf} and experimental searches, as well as being of intrinsic theoretical interest.  In this chapter we focus on the theoretical modelling of Hawking radiation from higher-dimensional black holes. We begin with a discussion of the quantum-field-theoretic derivation of Hawking radiation and its description using the Unruh vacuum state \\cite{Unruh:1976db}. We then briefly review some key features of black holes in brane world models. We describe the formalism for studying quantum fields on higher-dimensional Myers-Perry black holes \\cite{Myers:1986un}, which model black holes in an ADD brane-world \\cite{ArkaniHamed:1998rs,Antoniadis:1998ig,ArkaniHamed:1998nn}. We also present a selection of results on the properties of the Hawking radiation from these black holes. The literature on this subject is vast and so we cannot claim to do all aspects justice. The reviews \\cite{Kanti:2004nr,Casanova:2005id,Kanti:2008eq,Kanti:2012jh,Winstanley:2007hj} contain further discussion of results which space does not permit us to include. In the RS brane-world \\cite{Randall:1999ee,Randall:1999vf}, analysis of the Hawking radiation is more challenging because no exact metric describing a five-dimensional black hole localised on the brane is known in general - for a more detailed discussion of this topic, see the reviews \\cite{Gregory_Aegean, Kanti:2009sz, Tanahashi}. We close the chapter with a discussion of what is known about the Hawking radiation in this case. ",
        "conclusions": "\\label{sec:conc}  In this chapter we have reviewed some key aspects of our understanding of Hawking radiation from higher-dimensional black holes. One motivation for this study, apart from its intrinsic theoretical interest, is the exciting possibility of observing Hawking radiation from microscopic higher-dimensional black holes produced in high-energy collisions at either the LHC or in cosmic rays.  We began with a discussion of the quantum-field-theoretic foundations of Hawking radiation, from its original derivation for a quantum field on a dynamical space-time representing a black hole formed by gravitational collapse, to its modelling using the Unruh state on an eternal black hole geometry. After a brief review of classical black hole geometries in both the ADD and RS brane-worlds, most of this chapter has been devoted to the Hawking emission from these black holes.  For black holes in the ADD brane-world, for which we have an exact space-time metric, we gave an overview of the formalism used to describe the emission of massless particles both on the brane and in the bulk. We then reviewed a selection of results on the nature of the emission, including the important question of how much Hawking radiation escapes into the bulk and is therefore inaccessible to a brane-localised observer. Except for the emission of scalar- and vector-like graviton modes from a rotating black hole, our understanding of the emission of massless particles is essentially complete. We also briefly summarized some features of the Hawking radiation when other effects, such as the mass and charge of the emitted particles or a cosmological constant, are included.  In the RS brane-world, the lack of an exact classical metric for a black hole localized on the brane implies a scarcity of precise results on the nature of the Hawking emission from such black holes.  We have therefore just briefly outlined some of the work in the literature in this case.  Our review has revealed that Hawking radiation from higher-dimensional black holes is more complicated than from four-dimensional black holes, with many features depending on the number of extra dimensions and the angular momentum of the black hole. Many of the results we have outlined in this chapter have been incorporated into simulations of black hole events at the LHC \\cite{Dai:2007ki,Frost:2009cf}. Hawking radiation from black holes at the LHC characteristically involves large numbers of energetic particles, the details of the distribution and nature of the particles depending on the number of extra dimensions, the mass of the black hole and its angular momentum. Searches for these types of events have been made at the LHC, with no evidence to date for black holes \\cite{Aad:2011bw,Aad:2012ic,Aad:2013lna,Khachatryan:2010wx,Chatrchyan:2012taa,Chatrchyan:2013xva}. This non-observation of black hole events sets lower bounds on the energy scale of quantum gravity. With the LHC planned to run at higher energies and searches for high-energy black hole events in cosmic rays,  accurate modelling of Hawking radiation from higher-dimensional black holes continues to be essential for experimental probes of quantum gravity.  \\begin{acknowledgement} The work of E.W. is supported by the Lancaster-Manchester-Sheffield Consortium for Fundamental Physics under STFC Grant ST/J000418/1. The work of P.K. has been co-financed by the European Union (European Social Fund - ESF) and Greek national funds through the Operational Program ``Education and Lifelong Learning'' of the National Strategic Reference Framework (NSRF) - Research Funding Program: ``ARISTEIA. Investing in the society of knowledge through the European Social Fund''. Part of this work was supported by the COST Actions MP0905 ``Black Holes in a Violent Universe'' and MP1210 ``The String Theory Universe''. \\end{acknowledgement}"
    },
    "1402/1402.4575_arXiv.txt": {
        "abstract": "Among quadruples or higher multiplicity stars, only a few doubly eclipsing binary systems have been discovered. They are important targets to understand the formation and evolution of multiple stellar systems because we can obtain accurate stellar parameters from photometric and spectroscopic studies. We present the observational results of this kind of rare object 1SWASP J093010.78+533859.5, for which the doubly eclipsing feature had been detected previously from the SuperWASP photometric archive. Individual PSF photometry for two objects with a separation of about 1.9 arcsec was performed for the first time in this study. Our time-series photometric data confirms the finding of \\citet{lohr2013} that the bright object A is an Algol-type detached eclipsing binary and the fainter B is a W UMa-type contact eclipsing. Using the high-resolution optical spectra, we obtained well-defined radial velocity variations of system A. Furthermore, stationary spectral lines were detected that must have originated from a further, previously unrecognized stellar component. It was confirmed by the third object contribution from the light curve analysis. No spectral feature of the system B was detected, probably due to motion blur by long exposure time. We obtained the binary parameters and the absolute dimensions of the systems A and B from light curve synthesis with and without radial velocities, respectively. The primary and secondary components of system A have a spectral type of K1 and K5 main sequences, respectively. Two components of system B have nearly the same type of K3 main sequence. Light variations for both binaries are satisfactorily modeled by using two-spot models with one starspot on each component. We estimated the distances to systems A and B individually. Two systems may have similar distances of about 70 pc and seem to be gravitationally bound with a separation of about 130 AU. In conclusion, we suggest that 1SWASP J093010.78+533859.5 is a quintuple stellar system with a hierarchical structure of a triple system A(ab)c and a binary system B(ab). ",
        "introduction": "A stellar system consists of a small number of stars that are gravitationally bound and revolve each other around their center of mass. Since the stars in such a system are born nearly at the same time within a given space and evolve with interaction each other, they are important objects for understanding the stellar formation and dynamical evolution. The updated multiple star catalogue \\citep[MSC;][]{tokovinin1997} contains totally 1,359 stellar systems of multiplicity 3 to 7 in 2010  (at http://www.ctio.noao.edu/\\textasciitilde{}atokovin/). Among the most frequent triple stars, triply eclipsing hierarchical triple systems are very rare and fascinating to show peculiar light variations caused by the distant third component passing across the close eclipsing binary. To our knowledge, these tertiary eclipsing signals have been detected only in three systems, HD 181068 \\citep{derekas2011}, KOI-126 \\citep{carter2011}, and KIC 2856960 \\citep{armstrong2012,jwlee2013}, all from the $Kepler$ space telescope.  Another interesting object is a doubly eclipsing multiple system, which consists of at least two eclipsing binaries. We can investigate their physical characteristics in detail because the stellar mass and radius for each component are accurately and directly derived from both spectroscopic and photometric observations. They are also rare and only a handful of systems have been discovered (\\citealp[ADS 9537AB,][]{batten1965}; \\citealp[V994 Her,][]{leecu2008}; \\citealp[OGLE LMC-ECL-16549,][]{graczyk2011}; \\citealp[KIC 4247791,][]{lehmann2012}; \\citealp[CzeV343,][]{cagas2012}). On the other hand, \\citet{ofir2008} detected two eclipsing light variations with periods of about 3.58 and 5.37 days in the object OGLE J051343.14$-$691837.1 (= BI 108), which is located at the open cluster NGC 1881 in the Large Magellanic Cloud. He suggested that the object is a compact hierarchical system of two eclipsing binary pairs in a 3:2 resonance. Later, \\citet{kolaczkowski2013} performed spectroscopic observations and confirmed radial velocity variations by the longer period binary. However, unexpectedly no spectral feature by the shorter period binary was detected. Although its origin of variability may remain unclear, they suggested that the most plausible hypothesis is a doubly eclipsing quadruple system with the binary orbital period ratio of precise 3:2 by chance or by unknown dynamical reasons.  Recently, \\citet{lohr2013} reported the discovery of a new doubly eclipsing multiple system 1SWASP J093010.78+533859.5 (hereafter J093010) from the SuperWASP \\citep[Wide Angle Search for Planets,][]{pollacco2006} archive. The folded light curves with two periods ($\\sim$0.2277 and $\\sim$1.3055 days) show typical shapes of the contact and detached eclipsing binaries, respectively. The J093010 has been known to be a visual binary, TYC 3807-759-1 (hereafter J093010A) and TYC 3807-759-2 (hereafter J093010B), separated by 1.9 arcsec. Although the SuperWASP data could not resolve these two objects, \\citet{lohr2013} suggested that the light variations originated from a doubly eclipsing quadruple system. They calculated the separation between the two binaries to be 66.1 AU, by applying the distance of 35 pc estimated from the  photometric parallax by \\citet{agueros2009}, and derived their meta-orbital period to be $\\sim$400 years. It should be noted that the proper motion from the Tycho was measured at the photocenter of J093010 \\citep[][]{fabricius2002}, not for two individual objects. It may still be uncertain, therefore, whether the two objects are gravitationally bound.  In this study, we investigated the hierarchical structure of the multiple system J093010 and the physical properties of each component in detail from the follow-up spectroscopic and photometric observations. In Section 2, we describe the follow-up observations. The data analysis and modeling are represented in Sections 3 and 4, respectively. Finally, the summary and discussion are given in Section 5. ",
        "conclusions": "We performed photometric and spectroscopic follow-up observations for J093010 showing dual types of light curves. Our PSF photometry revealed the brighter object J093010A to be an Algol-type eclipsing binary and the fainter J093010B to be a W UMa-type eclipsing binary. Furthermore, the third component of J093010A was detected via spectroscopic observations and its light contribution was calculated to be about 30\\%. However, we could not detect any line in the spectra of J093010B probably due to relatively long exposure time to its orbital period. The RVs of J093010A were measured and its absolute dimensions were derived from the simultaneous analysis of light curves and RVs. The masses and radii of the components were found with an accuracy of $\\sim$2\\% and $\\sim$1\\%, respectively.  We compared our accurate physical parameters of J093010A with Dartmouth Stellar Evolution Program (DSEP) models \\citep{dotter2008}. First, several evolutionary tracks with different metallicities were compared to the position of the primary component in the HR diagram. In the upper panel of Figure \\ref{Fig6}, three evolutionary tracks of  M = 0.75 M$_\\odot$ are displayed on the HR diagram. The primary component is on the track of [Fe/H] = $-$0.25 (dashed line). Assuming the same metal abundances of two components, the secondary star (M = 0.644 M$_\\odot$) is located on the right-hand side of the track for [Fe/H] = $-$0.25 and M = 0.65 M$_\\odot$. We marked the positions of 9, 11, and 13 Gyr with filled circles in the panel. We also compared our parameters of J093010A with the empirical mass-radius relation from well-studied detached binaries (from DEBCat at http://www.astro.keele.ac.uk/jkt/debcat/). While the primary component is on the typical trend of main-sequence stars and on the isochrone of Age = 11 Gyr and [Fe/H] = $-$0.25, the secondary component has a larger radius than typical ones and is on quite above the isochrone (see lower panel of Figure \\ref{Fig6}). From the locations on the figure, the secondary component of J093010A has a relatively larger size and lower surface temperature than those expected from the model and empirical data. The main reason for these disagreements might be associated with the magnetic activities that caused the cool spot on the surface and H$\\alpha$ emissions \\citep{torres2013}. Chromospheric active stars showing H$\\alpha$ emission (see Figure \\ref{Fig3}) have a lower T$_{\\rm{eff}}$ and larger radius relative to both the theoretical isochrones and the non-active stars \\citep{stassun2012}. All these magnetic activities and changes of the light maximum are well-known characteristics of RS CVn type systems. The count rate in the X-ray band (0.1 - 2.4 keV) for J093010 is 0.433 counts s$^{-1}$ \\citep{voges1999} and the corresponding X-ray luminosity is calculated to be $\\sim$10$^{30}$ ergs s$^{-1}$, assuming the distance of 70 pc. This value is in the typical range for the RS CVn systems \\citep{dempsey1993}.  \\citet{lohr2013} suggested J093010 to be a quadruple system that consists of two binaries with a common proper motion. As noted in the Tycho Double Star Catalogue \\citep{fabricius2002}, it is the photocentre proper motion for a close double which was neither resolved in Tycho-2 nor measured individually. We calculated the distances to each binary to be 66 $\\pm$ 7 and 77 $\\pm$ 9 pc, respectively. The value for J093010B was estimated from an empirical relation between the temperature and stellar mass, and hence the error may be larger than the presented one. Two eclipsing binaries have almost same inclinations, similar masses and spectral types, which seem to be located at the same distance and to be gravitationally bound. The distance to the third component of J093010A, assuming a main-sequence star, is also comparable to those of the two binaries. Moreover, the system velocity of the Algol-type binary is almost consistent with the RVs of the stationary component. This means that the third body is possibly bound to the eclipsing pair. In conclusion, we suggest that J093010 is a possible quintuple system with a hierarchical structure of a triple A(ab)c and a binary B(ab). Assuming that the distance to the system is 70 pc and the mass of the third component of J093010A is 0.7 M$_\\odot$, the separation and meta-orbital period between the two systems (J093010A and B) are calculated to be $\\sim$130 AU and $\\sim$850 years, respectively. While the RVs for the third component of J093010A seem to have slightly decreased by about 1 km s$^{-1}$ during our observation interval of $\\sim$100 days, the corresponding increment ($\\sim$0.5 km s$^{-1}$) of system velocities for J093010A(ab) was too small to measure. High-precision and long-term observations are required to detect the variations. In addition, the spectroscopic observations for J093010B will be helpful to obtain more accurate physical parameters and distance to the binary.       \\newpage"
    },
    "1402/1402.2059_arXiv.txt": {
        "abstract": "We extend previous classifications of inflationary models by means of their behaviour at large-$N$, where $N$ is the number of e-foldings. In addition to the perturbative $1/N$ case, whose slow-roll parameters fall off as powers of $1/N$, we introduce the constant, non-perturbative and logarithmic classes. This covers the large majority of inflationary models. Furthermore, we calculate the running of the spectral tilt for all these classes. Remarkably, we find that the tilt's runnings essentially cluster around the permil level. We comment on the implications for future experiments. ",
        "introduction": "After measuring the cosmic microwave background (CMB) with unprecedented precision, Planck has presented us with a pristine image of the early Universe. This has resulted in strong constraints on the inflationary phase. In particular, the spectral index has been measured to be \\cite{Ade:2013} \\begin{align} n_s & = 0.9603 \\pm 0.0073 \\,, \\label{ns} \\end{align} although subsequent analyses have argued for a slightly higher value \\cite{Spergel:2013}. Similarly, the running of the spectral index has been constrained with some precision: \\begin{align} \\alpha_s = - 0.0134 \\pm 0.0090 \\,. \\end{align} On the other hand, the BICEP2 collaboration~\\cite{BICEP2} has recently announced the detection of the primordial gravitational wave background produced during inflation, with an amplitude responsible for a tensor-to-scalar ratio \\begin{align} r = 0.20^{\\,+0.07}_{\\,-0.05} \\,, \\end{align} with $r=0$ ruled out at more than 5$\\sigma$. This bound is in slight tension with the reported Planck upper bound for a 6-parameter $\\Lambda$CDM model, $r < 0.11$ at 95\\% c.l., although foreground subtraction could decrease to some extent  the reported value. If confirmed, such detection would lead to important consequences. In particular, many inflationary scenarios are ruled out by these bounds. Indeed, Planck and BICEP2 together seem to prefer large field models of inflation~\\cite{Martin:2013a, Martin:2013b}, with predictions different from the plateau potentials of Starobinsky-like models~\\cite{Starobinsky:1982}, preferred before BICEP2 results.  In this paper we aim to understand the values quoted above in terms of the number of e-foldings $N$ between the horizon exit of the quantum modes that have resulted in the CMB anisotropies, and the end of inflation. As will be discussed in more detail below, this approach allows for a classification of different inflationary models in terms of universality classes with distinct large-$N$ behavior. It builds on preceeding works \\cite{Mukhanov:2013, Roest:2013}, where the deviation from a Harrison-Zeldovich scale-invariant spectrum with $n_s = 1$ was attributed to $1/N$ effects. We will extend this with different leading terms in the $1/N$ expansion. Moreover, we will also investigate the running of the spectral index as a function of $N$. The running is the second-order term in the expansion of the scale-dependence of the power spectrum, and hence is natural to consider it in this framework as well. Finally, the values of this running is considered for a wide range of models, and shown to cluster in a very limited range around a central value of $\\log_{10}|\\alpha_s| = - 3.2$.  The outline of the paper is as follows. We introduce the $N$-formalism in section 2, where all slow-roll parameters are expressed in terms of $N$. Section 3 contains the different universality classes of large-$N$ behaviour. The running of the spectral index is discussed in section 4. Finally, we conclude in section 5. In the Appendix we consider also models of $k$-inflation as a separate class of its own. ",
        "conclusions": "In this paper we have introduced the $N$-formalism, which can be thought of as the background complement to the $\\delta N$-formalism and where all inflationary observables are expressed in terms of the number of e-foldings. This is particularly useful given the generic requirement that $N$ should exceed $50$ and suggests a large-$N$ expansion.  We have demonstrated that the majority of models  allow for such an expansion in $1/N$; the leading terms for $n_s$ are either constant, perturbative, non-perturbative or logarithmic. A quick overview of the properties of these classes can be found in table 1. Moreover, we have shown that subleading terms are generically irrelevant (in addition to ambiguously defined): the leading approximation at large-$N$ agrees excellently with the exact expression. The only exception to this behaviour are the non-perturbative examples of natural and quadratic hilltop. In these models, the effective expansion is in terms of $cN$, which does not allow for a leading approximation when $c$ is sufficiently small.  \\begin{table}[ht] \\begin{center} \\begin{tabular}{||c||c||c||} \\hline Class & $V(\\phi)$ & $\\epsilon(N)$  \\\\ \\hline \\hline Constant & $e^{\\phi}$ & constant \\\\ \\hline Perturbative $p=1$ & $\\phi^n$ with $n>0$ & $1/N$ \\\\ \\hline Perturbative $1<p<2$ & $1- \\phi^n$ with $n<0$ & $1/N^p$ \\\\ \\hline Perturbative $p=2$ & $1-e^{-\\phi}$ & $1/N^2$ \\\\ \\hline Perturbative $p>2$ & $1- \\phi^n$ with $n>0$ & $1/N^p$ \\\\ \\hline Non-perturbative & $1-\\phi^2$ & $e^{-N}$ \\\\ \\hline Logarithmic & $1-\\phi e^{-\\phi}$ & $\\ln(N)/N$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{A sketch of the various classes and their corresponding asymptotic potentials and equations of state at large-$N$. More details can be found in section 3.} \\end{table}  We should also mention that there are inflationary models that do not allow for a large-$N$ expansion. The chief example is supergravity hybrid inflation~\\cite{LindeRiotto} with a potential \\be V(\\phi) = \\mu^4\\left( 1 + {\\kappa^2\\over8\\pi^2}\\ln\\frac{\\phi}{\\phi_c} + \\frac{\\phi^4}{8}\\right) \\,. \\ee The resulting equation of state parameter reads \\be \\epsilon = \\frac{\\kappa^3}{16\\pi^3\\sin\\frac{\\kappa N}{\\pi}(1+\\cos\\frac{\\kappa N}{\\pi})}\\,. \\ee Due to the trigonometric functions, this parameter does not have a definite large-$N$ behavior and hence defies the proposed classification.  Building on the $N$-dependence of these universality classes, we have investigated the values of the running for the different inflationary models. Using the Planck constraints for the spectral index as input, it turns out that essentially all models give a similar prediction for the running of $n_s$ as a function of scale, centered around $\\log_{10}|\\alpha_s| = - 3.2$ and with a range of only half a decade. This remarkable feature seems to have gone unnoticed so far. As we have pointed out, one can gain a first understanding of the limited range of $\\alpha_s$ from the preference of Planck for inflationary models with correct spectral indices for both $N=50$ as well as $N=60$."
    },
    "1402/1402.5954_arXiv.txt": {
        "abstract": "We have carried out a search for optically visible post-AGB candidates in the Small Magellanic Cloud (SMC).  Firstly, we used mid-IR observations from the Spitzer Space Telescope to select optically visible candidates with excess mid-IR flux and then we obtained low-resolution optical spectra for 801 of the candidates.  After removing poor quality spectra and contaminants such as M-stars, C-stars, planetary nebulae, quasi-stellar objects and background galaxies, we ended up with a final sample of 63 high probability post-AGB/RGB candidates of A $-$ F spectral type. From the spectral observations, we estimated the stellar parameters: effective temperature ($T_{\\rm eff}$), surface gravity ($\\log g$), and metallicity ([Fe/H]). We also estimated the reddening and deduced the luminosity using the stellar parameters combined with photometry. For the post-AGB/RGB candidates, we found that the metallicity distribution peaks at [Fe/H] $\\approx$ -1.00 dex. Based on a luminosity criterion, 42 of these 63 sources were classified as post-RGB candidates and the remaining 21 as post-AGB candidates. From the spectral energy distributions we were able to infer that  6 of the 63 post-AGB/RGB candidates have a surrounding circumstellar shell suggesting that they are single stars, while 27 of the post-AGB/RGB candidates have a surrounding disc, suggesting that they lie in binary systems. For the remaining 30 post-AGB/RGB candidates the nature of the circumstellar environment was unclear. Variability is displayed by 38 of the 63 post-AGB/RGB candidates with the most common variability types being the Population II Cepheids (including RV-Tauri stars) and semi-regular variables. This study has also revealed a new RV Tauri star in the SMC, J005107.19-734133.3, which shows signs of $s$-process enrichment. From the numbers of post-AGB/RGB stars in the SMC, we were able to estimate evolutionary rates. We find that the number of post-AGB and post-RGB candidates that we have identified are in good agreement with the stellar evolution models with some mass loss in the post-AGB phase and a small amount of re-accretion in the lower luminosity post-RGB phase. This study also resulted in a new sample of 40 young stellar objects (YSOs) of A $-$ F spectral type. The 40 YSO candidates for which we could estimate stellar parameters are luminous and of high mass ($\\sim3-10 \\Msun$). They lie on the cool side of the usually adopted birthline in the HR-diagram. This line separates visually obscured protostars from optically visible pre-main sequence stars, meaning that our YSO candidates have become optically visible in the region of the HR-diagram usually reserved for obscured protostars. Additionally, we also identified a group of 63 objects whose spectra are dominated by emission lines and in some cases, a significant UV continuum. These objects are very likely to be either hot post-AGB/RGB candidates or luminous YSOs. ",
        "introduction": "\\label{intro} The post-asymptotic giant branch (post-AGB) phase of evolution is a transient phase between the asymptotic giant branch (AGB) and planetary nebula (PNe) phases of stellar evolution. A super-wind with mass-loss rates up to $10^{-4}$M$_{\\odot}$\\,yr$^{-1}$ generally terminates the AGB phase, reducing the hydrogen-rich envelope to small values. Subsequently, the radius of the central star decreases and within about $10^{2}-10^{4}$yrs, the star evolves to higher temperatures (from 3$\\times10^{3}$K on the AGB to $\\sim$3$\\times10^{4}$K where PN formation begins) with a constant luminosity \\citep[e.g.][]{schoenberner83,vw93,vw94}. The ejected circumstellar matter expands and cools during the post-AGB phase, resulting in stars with a large mid-IR excess \\citep[see][for reviews]{kwok93,vanwinckel03,habing03-bk}. Post-AGB stars emit radiation over a broad spectral range owing to a combination of high temperatures in  the photosphere and low temperatures in circumstellar dust. This enables the simultaneous study of the stellar photosphere and the circumstellar environment, with the central star emitting in the ultra-violet(UV), optical and near-infrared (IR), and the cool circumstellar environment radiating in the near- and mid-IR. As the post-AGB star evolves to higher temperatures, the ejected circumstellar material gets ionised to form a PNe, after which the central star declines in luminosity and cools to become a white dwarf star. These phenomenon mark the end of stellar evolution for low-to intermediate-mass single stars.  For stars in binary systems, a different mechanism can terminate the red giant evolution. The large expansions that occurs when a star becomes a red giant can cause the primary star to over-fill its Roche lobe. Theory suggests that Roche lobe overflow during the red giant phase will lead to run-away mass transfer on a dynamical or thermal timescale, until all that is left of the mass-losing star is the He or C/O core of the red giant, orbiting in a binary system: alternatively, complete merging may occur \\citep[e.g.][]{paczynski72,iben96,han95a}. This process occurs for binaries with periods on the main sequence in the approximate range 20 $-$ 1000 days, from low on the red giant branch (RGB) to the tip of the AGB. The outcome of these systems are as of yet difficult to predict as many binary interaction process are poorly understood. A significant fraction of the ejected matter may end up in a circumbinary disc of dust and gas, and inside the disc is a binary system containing a post-AGB star or a post-RGB star. The distinction between these two cases is as follows: if the primary star fills its Roche lobe before reaching the RGB tip\\footnote{According to observational studies in the Magellanic Clouds \\citep[e.g.,][]{frogel83,wood99-pr,cioni99} and evolutionary tracks of \\citet{bertelli08} corresponding to LMC and SMC metallicities, the bolometric luminosity of the RGB tip is found to be close to $M_{\\rm bol} -3.6$ mag.}, it will be referred to as a post-RGB star while if the primary star successfully evolves past the RGB tip on its second ascent of the giant branch it will be referred to as a post-AGB star.  One of the challenges in the study of the post-AGB phase of evolution is the identification of post-AGB objects as they have a very short lifetime. Since they have dusty circumstellar envelopes, the detection of cold circumstellar dust using mid-IR photometry is an efficient method to select and study them. The first extensive search for these objects was initiated in the mid-80's using results from the Infrared Astronomical Satellite (IRAS). The large scale mid-IR IRAS mission enabled the identification of post-AGB stars in the Galaxy \\citep{kwok93}. The Toru$\\acute{\\rm n}$ catalogue \\citep{szczerba07} for Galactic post-AGB stars lists around 391 very likely post-AGB objects. The Galactic sample of optically visible post-AGB objects has revealed two highly distinct populations: one with cold, detached, expanding dust shells \\citep[these probably arise from single stars and produce 'shell' or 'outflow sources',][]{vanwinckel03}, and another with hot dust and circumstellar discs \\citep[these arise from binary stars and are called 'disc sources',][]{deruyter06,vanwinckel07,gielen09,vanwinckel09,dermine12}. This is as expected from the single and binary star evolution scenarios described above.  So far, in the Galaxy, the luminosities (and hence initial masses) of the diverse group of post-AGB stars are badly affected by their unknown distances, making it difficult to use the observational characteristics of these interesting objects to throw light on the poorly-understood late stages of stellar evolution.  The Magellanic clouds are well suited for the identification and observation of post-AGB stars. The well-constrained distances to these extragalactic systems mean that distance-dependent parameters such as luminosities can be determined accurately. The Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) are both very suitable environments to locate individual post-AGB objects and study their evolution as a function of initial mass and metallicity.  In the SMC, only 5 possible post-AGB candidates have been identified previously. One IRAS source (IRAS 00350-7436) is believed to be a post-AGB star \\citep{whitelock89}. \\citet{kuzinskas00} identified a near-IR object in the field of the SMC cluster NGC\\,330 using ISO (Infrared Space Observatory) observations. This object was classified as a likely post-AGB star in a proto-planetary nebula stage, or a Be-supergiant. Using low-resolution mid-IR spectra from the Spitzer Space Telescope (SST), \\citet{kraemer06} identified a possible isolated post-AGB star (MSX SMC 029) in the SMC. \\citet{volk11} identified two 21 micron sources in the SMC (J004441.04-732136.44 and J010546.40-714705.27). For the Galaxy the unidentified 21 micron feature only occurs around c-rich post-AGB stars. J004441.04-732136.44 was also identified in a preliminary version of this study by \\citet{wood11}, who examined some of the brighter sources presented here and found a number of post-AGB candidates. \\citet{desmedt12} carried out a detailed chemical abundance analysis on J004441.04-732136.44 and found it to be indeed C-rich and an extremely \\emph{s}-process enriched shell-source.  In the LMC, post-AGB candidates have been identified by \\citet{vanaarle11} and references therein. \\citet{vanaarle11} constructed a catalogue of 1337 optically visible post-AGB candidates in the LMC based on mid-IR selection criteria and examination of spectral energy distributions (SEDs). They also carried out a spectroscopic examination of 105 candidates which resulted in 70 likely post-AGB candidates in the LMC.  To fully understand the post-AGB/RGB population in the Magellanic Clouds and to provide a better sample of post-AGB/RGB stars whose properties can be compared to models, a larger sample of these objects is definitely required. We have therefore carried out an extensive search for post-AGB/RGB stars in the SMC, to complement the survey of \\citet{vanaarle11} in the LMC. In this paper, we present the results of our survey of post-AGB/RGB stars in the SMC. This includes a low-resolution optical spectral survey of the post-AGB candidates.\\footnote{We have also carried out a low-resolution optical spectroscopic examination for the sample of the LMC post-AGB candidates identified by \\citet{vanaarle11}. The results will be presented in a following publication (Kamath et al., in preperation).}  The structure of this paper is as follows: In Section~\\ref{sampleselection} we present the selection criteria used to obtain an initial sample of post-AGB/RGB candidates in the SMC. We also present the photometric data, covering the UV, optical and IR wavelengths and use it to compute the SEDs for all the sources in our sample. Optical spectra along with the data reduction procedure and a preliminary spectral classification, are presented in Section~\\ref{specobv}. In Section~\\ref{STP}\\,$-$\\,~\\ref{ms}, we describe the method used to estimate the stellar parameters (\\teff\\,, \\logg\\,, and [Fe/H]) from the spectra, and the reddening from SED fitting. Subsequently, in Section~\\ref{classification} we present the catalogues of the post-AGB/RGB and YSO candidates. In Sections~\\ref{sedanalysis}\\,$-$~\\ref{lcanalysis}, we analyse different characteristics of the two populations by examining the SEDs, the optical spectra and the light curves of the individual candidates. In Section~\\ref{completeness} we present the completeness of the survey and in Section~\\ref{evolrate} we empirically estimate the evolutionary rates of the transient post-AGB/RGB phase. Finally, we conclude with a summary of our survey. ",
        "conclusions": ""
    },
    "1402/1402.5353_arXiv.txt": {
        "abstract": "% In the context of high-frequency quasi-periodic oscillation (HF QPOs) we further explore the appearance of an observable signal generated by hot spots moving along quasi-elliptic trajectories close to the innermost stable circular orbit in the Schwarzschild spacetime. The aim of our investigation is to reveal whether observable characteristics of the Fourier power-spectral density can help us to distinguish between the two competing models, namely, the idea of bright spots orbiting on the surface of an accretion torus versus the scenario of intrinsic oscillations of the torus itself. We take the capabilities of the present observatories (represented by the Rossi X-ray Timing Explorer, RXTE) into account, and we also consider the proposed future instruments (represented here by the Large Observatory for X-ray Timing, LOFT). ",
        "introduction": "\\renewcommand*{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext{Talk presented at 10$^{\\rm th}$ INTEGRAL/BART Workshop in Karlovy Vary (Czech Republic, 22-25 April 2013).} Low-mass X-ray binaries (LMXBs) represent a particular example of astronomical sources in which puzzling quasi-periodic modulation of the observed signal can develop and reach very high (kilohertz) frequencies \\cite{Kli:2006:CompStelX-Ray:}. Orbiting inhomogeneities, a.k.a.\\ spots residing on the surface of an accretion disk, have been introduced as a tentative explanation for persistant high-frequency oscillations (HF QPOs) observed in X-rays from accreting black holes and neutron stars in LMXBs. These QPOs have been discussed, in particular, in a series of papers \\citep{kar:1999,mor-ste:1999,ste-vie:1998a,ste-vie:1998b,ste-vie:1999,ste-vie:2002} that explain QPOs as a direct manifestation of modes of relativistic epicyclic motion of blobs at various radii $r$ in the inner parts of the accretion disc. Within the model, the two observed HF QPO peaks arise due to the Keplerian and periastron precession of the relativistic orbits.  In a similar manner, the authors of refs.\\ \\cite{cad-etal:2008,ger-etal:2009,kos-etal:2009} introduced a concept according to which the QPOs are generated by a tidal disruption of large accreting inhomogeneities. Recently, \\cite{ger:2012} argued that the behaviour of  the azimuth phase $\\phi(t)$ for non-closed quasi-elliptic orbits in the curved spacetime can be responsible for the observed pairs of HF QPOs.  In this contribution we investigate the behaviour of the observable signal produced by radiating small circular hotspots. We discuss the detectability of the produced signal propagated from the strongly curved spacetime region, and in particular we concentrate of the properties of Fourier power-spectral density (PDS), as predicted in different scenarios. It has been known that PDS reflects the physical properties of the source, however, it does not provide a complete description. We thus explore a possibility to discriminate via PDS between the two specific models that are thought to provide a promising scheme for HF QPOs in black hole sources. In fact, according to the Wiener-Khinchin theorem, the power spectral function $S(\\omega)$ can be calculated from the autocovariance function $C(t)$ as \\cite{bendat:2000} \\begin{equation} S(\\omega)={\\cal F}[C](\\omega). \\end{equation} It can be seen that only the Gaussian processes are completely determined by their power spectra. Nonetheless, it is interesting to investigate the predicted profiles of the power spectrum for well-defined processes and to reveal specific features, their differences and similarities.  In our discussion we consider the capabilities of the present X-ray observatories represented by Rossi X-ray Timing Explorer (RXTE), as well as the proposed future instruments represented by the Large Observatory for X-ray Timing (LOFT). We also compare the signal produced by spots to the signal obtained from another specific kind of simulations assuming axisymmetric epicyclic disc-oscillation modes. Our present paper is based on a recent work \\citep{bakala:2013} where further details can be found. In particular, in the cited paper we explore the width of the QPO peaks with 3:2 frequency ratio that has been reported to occur in a number of sources. ",
        "conclusions": "We can identify the signatures of the spot motion mostly with the harmonic content of the observable signal. For large inclination angles, the LOFT observations could easily reveal the Keplerian frequency of the spot together with its first and second harmonics when the strongest (but weak) single signal is around the limits of the RXTE detectability. Nevertheless, radial epicyclic frequency could be also found providing that the inclination is small. In our analysis we have paid attention to the timing signatures of the motion of small circular spots radiating isotropically from the slightly eccentric geodesic orbits. The case of highly eccentric orbits and/or spot having large azimuthal shear will be presented elsewhere.  We studied the comparison between spot and torus scenarios serving as  fiducial representations of two very specific kinematic models. Obviously, any general validity of this discussion is limited. For instance, a consideration of resonance driven effects or the role of torus geometrical thickness could give rise to some harmonic content in the signal from the oscillating tori. Despite uncertainties, the elaborated comparison indicates clearly that the increased sensitivity of the proposed LOFT mission can be crucial for resolving the QPO nature. We refer the reader to ref.\\ \\cite{bakala:2013} for additional details.    \\thanks We acknowledge the project CZ.1.07/2.3.00/20.0071 Synergy supporting the international collaboration of Institute of Physics in Opava, the Czech Science Foundation project GA\\v{C}R 13-00070J in Prague (VK), GA\\v{C}R~209/12/P740 (GT, E\\v{S}), and the European Union 7th Framework Programme No. 312789 `StrongGravity' (MD). We also acknowledge the Polish grant NCN UMO-2011/01/B/ST9/05439 and the Swedish VR grant (MAA). Astronomical Institute is supported by the research program RVO:67985815."
    },
    "1402/1402.7026_arXiv.txt": {
        "abstract": "{We consider the Lagrangian of a vector field with derivative self-interactions with a priori arbitrary coefficients. Starting with a flat space-time we show that  for a special choice of the coefficients of the self-interactions the ghost-like pathologies disappear. This constitutes the Galileon-type generalization of the Proca action with only three propagating physical degrees of freedom. The longitudinal mode of the vector field is associated to the usual Galileon interactions for a specific choice of the overall functions. In difference to a scalar Galileon theory, the generalized Proca field has more free parameters. We then extend this analysis to a curved background. The resulting theory is the Horndeski Proca action with second order equations of motion on curved space-times. } ",
        "introduction": "Motivated by the work from de Rham and Gabadadze for the generalization of the Fierz-Pauli action for a massive graviton \\cite{deRham:2010ik}, we investigate here the generalization of the Proca action for a massive vector field with derivative self-interactions. We will be addressing the natural question of what is the Lagrangian for a self-interacting vector field with second order equations of motion yielding three propagating physical degrees of freedom. We will call them the \"vector Galileons\", since they contain derivative self-interactions for the vector field and the longitudinal mode corresponds to a Galileon \\cite{Nicolis:2008in}.\\\\ In the standard relativistic quantum field theory we describe particles with local covariant field operators like scalars, vectors, tensors..etc. The finite-dimensional representation of the Lorentz group dictates to us the number of propagating degrees of freedom. For a massless spin-1 field the theory needs to have the gauge symmetry in order to have the Lorentz invariance manifestly built in. The theory describes then a massless spin-1 field with two propagating degrees of freedom $h=\\pm 1$. On the other hand, for a massive spin-1 field we have $(2s+1)$, i.e three, propagating physical degrees of freedom. The Proca action is the theory describing  a massive vector field, which propagates the corresponding three polarizations (two transverse plus one longitudinal) . The mass term breaks explicitly the $U(1)$ gauge invariance such that the longitudinal mode propagates as well. However, the zero component of the vector field does not propagate. So therefore it is a natural question to investigate also the existence of derivative interactions for the vector field with still only three propagating degrees of freedom. This is exactly what we aim in this paper: we want to find the generalization of the Proca action for a massive vector field with derivative self-interactions. The standard Proca action is given by \\begin{equation} \\mathcal S_{\\rm Proca}  = \\int \\mathrm{d}^4x \\left[-\\frac14 F_{\\mu\\nu}^2 -\\frac12 m^2 A^2 \\right]  \\label{ProcaAction} \\end{equation} where $F_{\\mu\\nu}=\\partial_\\mu A_\\nu -\\partial_\\nu A_\\mu$. In this theory, the temporal component of the vector field does not propagate and generates a primary constraint. The consistency condition of this primary constraint generates a secondary constraint, whose Poisson bracket with the primary constraint is proportional to the mass so that only in the massless case it corresponds to a first class constraint generating a gauge symmetry. Gauge invariance is just a redundancy in the description of massless particles describing the same physical state. Therefore, we can use the Stueckelberg trick to restore the gauge invariance. In the case of the standard Proca action we can restore the gauge invariance by adding an additional scalar field via $A_\\mu \\to A_\\mu + \\partial_\\mu \\pi$. This trick does not change the number of propagating physical degrees of freedom. We add one additional scalar degree of freedom but we restore the gauge invariance which guaranties the existence of only two physical degrees of freedom for the vector field, making it in total three physical degrees of freedom. When we now take the mass going zero limit  $m_\\to0$ the Lagrangian results in a theory of a massless scalar field completely decoupled from a massless vector field for a conserved source. This is the reason why there is no vDVZ discontinuity in the case of the Proca field for conserved sources. This is different for the massive graviton. There, the helicity-0 degree of freedom does not decouple and gives rise to an additional fifth force which has to be screened via Vainshtein mechanism. In the case of the vector field, we do not need any Vainshtein mechanism since there is no observational difference between a massless and massive vector fields for conserved sources. In the generalized Proca action that we construct here, the longitudinal mode of the vector field has exactly the same interactions as a Galileon scalar field for a specific choice of the overall functions. The Galileon theory is an important class of infra-red modifications of general relativity . The Galileon interactions were introduced as a natural extension of the decoupling limit of the DGP model \\cite{Nicolis:2008in}. It is constructed as an effective field theory for a scalar field by the restriction of the invariance under internal Galilean and shift transformations and second order equations of motion. This effective action is local and contains higher order derivatives. Nevertheless, these interactions come in a very specific way such that they only give rise to second order equations of motion. The allowed interactions for the Galileon were originally determined order by order by writing down all the possible contractions for the derivative scalar field interactions and finding the proper coefficients giving rise to second order equations of motion. However, de Rham and Tolley could construct an unified class of four dimensional effective theories starting from a higher dimensional setup and show that these effective theories reproduce successfully all the interaction terms of the Galileon in the non-relativistic limit \\cite{deRham:2010eu}. In a similar way we wonder whether or not one could construct the generalized Proca action that we are proposing here from a higher dimensional set-up which we will investigate in a future work. Naively, we would think, that, starting from a higher dimensional set-up with manifestly covariant Lovelock invariants, one would only construct terms which are gauge invariant after dimensional reduction. There is only one possible non-minimal interaction which fulfills this requirement, namely the contraction of two field strength tensors with the dual Riemann tensor. Therefore, this specific interaction with gauge invariance could be easily constructed from a higher dimensional set-up. However, it would be worth to study, if the other not gauge invariant non-minimal couplings could be constructed by dimensional reduction.   The Galileon interactions present a subclass of Horndeski interactions which describe scalar-tensor interactions with at most second order equations of motion on curved backgrounds \\cite{Horndeski:1974wa}. Interestingly, a subclass of Horndeski scalar-tensor interactions \\cite{deRham:2011by} can also be constructed by covariantizing the decoupling limit of massive gravity \\cite{deRham:2010kj}. In the literature there has been some attempts to find a theory for vector fields which is equivalent to scalar Galileons, i.e. to find the  \"vector Galileons\" besides the Maxwell kinetic term with second order equations of motion on flat space-times \\cite{Deffayet:2013tca}. There, the authors were interested in derivative self-interactions for the vector field with gauge symmetry yielding only two propagating degrees of freedom. They concluded that the Maxwell kinetic term is the only allowed interaction and wrote a no-go theorem for generalized vector Galileons \\footnote{A similar no-go theorem has been also studied in \\cite{deRham:2013tfa} in the context of massive graviton}. However, on curved background there only exists one term respecting the gauge symmetry, which is given by the non-minimal coupling between the field strength tensor and the dual Riemann tensor \\cite{Horndeski:1976gi, Jimenez:2013qsa}. However, if one gives up on the gauge invariance meaning that we allow for terms which are not invariant under $A_\\mu \\to A_\\mu + \\partial_\\mu\\theta$ then one can indeed construct \"vector Galileons\" on flat-spacetimes or Horndeski vector interactions on curved backgrounds giving rise to three propagating physical degrees of freedom with second order equations of motion. We will illustrate this in this work. ",
        "conclusions": "In this paper we have constructed the generalized Proca action for a vector field with derivative self-interactions with only three propagating degrees of freedom. We started our analysis with the case of a flat Minkowksi spacetime. We successfully showed that for appropriate choices of the coefficients of the derivative self-interactions that generalize the Proca action, one can construct a consistent and local theory of massive vector field without the presence of ghost-like instabilities. The resulting theory is simple and constitutes four Lagrangians for the self-interactions of the vector field. We were able to show that the constrained coefficients yield the necessary propagating constraint in order to remove the unphysical degree of freedom. These are the \"vector Galileons\" with three propagating degrees of freedom. At each order the Lagrangian has an overall function which depends on $A^2$ and the function of the quadratic Lagrangian can also depend on all the possible terms invariant under $U(1)$ symmetry like for instance $F^2$ and $FF^*$..etc. Similarly this function can also depend on any contractions between the vector field and the field strength tensor $A_\\mu A_\\nu F^{\\mu\\rho}F_\\rho^\\nu$ which does not contain any time derivative applied on the temporal component of the vector field. The dependence of the function $f_2$ on the gauge invariant terms or terms in which the zero component of the vector field does not have any dynamics, do not alter the number of propagating degrees of freedom. We have also shown, that these interactions have more free parameters than the corresponding scalar Galileon interactions. We then generalized our results to the case of curved space-time and obtained the corresponding \"Horndeski vector\" interactions. In the very latest stage of this work we became aware that a similar idea has been explored in \\cite{Tasinato:2014eka} even though the interactions we constructed here are more general."
    },
    "1402/1402.2329_arXiv.txt": {
        "abstract": "{} { In this work, we perform a complete census of molecular ions with an abundance larger than $\\sim 10^{-10}$ in the protostellar shock L1157-B1. This allows us to study the ionization structure and chemistry of the shock.} {An unbiased high-sensitivity survey of L1157-B1 performed with the IRAM-30m and {\\it Herschel}/HIFI as part of the CHESS and ASAI large programs allows searching for molecular ions emission. Then, by means of a radiative transfer code in the Large Velocity Gradient approximation, the gas physical conditions and fractional abundances of molecular ions are derived. The latter are compared with estimates of steady-state abundances in the cloud and their evolution in the shock calculated with the chemical model Astrochem. } {We detect emission from \\hcop, \\hccop, \\nnhp, \\hcsp, and, for the first time in a shock, from \\hocop, and \\sop. The bulk of the emission peaks at blueshifted velocity, $\\sim 0.5-3$~\\kms\\ with respect to systemic, has a width of $\\sim 4-8$~\\kms\\ and is associated with the outflow cavities (T$_{\\rm kin}\\sim20-70$~K,  \\denshh~$\\sim10^5$~\\cmc). A high velocity component, up to $-40$~\\kms, associated with the primary jet is detected in the \\hcop~1--0 line. Observed \\hcop\\ and \\nnhp\\ abundances (\\xhcop~$\\sim0.7-3~10^{-8}$, \\xnnhp~$\\sim0.4-8~10^{-9}$) are in agreement with steady-state abundances in the cloud and with their evolution in the compressed and heated gas in the shock for cosmic rays ionization rate $\\zeta=3~10^{-16}$~s$^{-1}$. % \\hocop, \\sop, and \\hcsp\\ observed abundances (\\xhocop~$\\sim10^{-9}$, \\xsop~$\\sim8~10^{-10}$, \\xhcsp~$\\sim3-7~10^{-10}$), instead, are 1--2 orders of magnitude larger than predicted in the cloud; on the other hand they are strongly enhanced on timescales shorter than the shock age ($\\sim$2000 years) if \\coo, S or H$_2$S, and OCS are sputtered off the dust grains in the shock. } {The performed analysis indicates that \\hcop\\ and \\nnhp\\ are a fossil record of pre-shock gas in the outflow cavity, while \\hocop, \\sop, and \\hcsp\\ are effective shock tracers and can be used to infer the amount of \\coo\\ and sulphur-bearing species released from dust mantles in the shock. The observed \\hcsp\\ (and CS) abundance indicates that OCS should be one of the main sulphur carrier on grain mantles. However, the OCS abundance required to fit the observations is 1--2 orders of magnitude larger than observed. Laboratory experiments are required to measure the reactions rates involving these species and to fully understand the chemistry of sulphur-bearing species.} ",
        "introduction": "The gravitational infall from which a new star is formed is accompanied by the ejection of highly supersonic jets. The shocks produced when the ejected material impacts on the high-density surrounding medium rapidly heat and compress the gas triggering several microscopic processes such as molecular dissociation, gas ionization, endothermic chemical reactions, ice sublimation, and dust grain disruption. It is well know that these processes produce a significant enhancement of the abundance of molecules which are evaporated or sputtered from dust grains, such as H$_2$CO, CH$_3$OH, NH$_3$, SiO  \\citep[see, e.g., ][]{bachiller97}.  Molecular ions are commonly observed in dense prestellar/protostellar regions \\citep[e.g., ][]{caselli98,caselli02c,hogerheijde98} and their chemistry have been studied by means of gas-phase chemical models \\citep[e.g., ][]{herbst86,turner99}. In contrast, there are very few observations of molecular ions in protostellar shocks. \\hcop\\ was detected in a number of outflows and in some cases show a 'butterfly' morphology suggesting excitation in the outflow cavities \\citep[e.g., ][]{hogerheijde98,girart99,girart02,lee07,tafalla10,tappe12}. Besides \\hcop, emission in \\hcsp\\ and very recently in \\nnhp\\ lines was detected in the B1 shock along the chemically rich L1157 outflow \\citep{bachiller97,yamaguchi12,codella13}.  Due to the lack of observations, the chemical evolution of molecular ions in shocks has been poorly investigated. On the one hand, molecular ions are expected to be rapidly destroyed by electronic recombination in shocks; on the other hand, they may be enhanced due to the sputtering of dust grains mantles, which injects molecules in the gas phase where they undergo reactions leading to ions \\citep[e.g., ][]{neufeld89a,viti02,rawlings04}. Among the  molecular ions previously observed in shocks, not all of them are due to changes caused by the shock itself. \\citet{codella13} show that the \\nnhp\\ abundance is expected to decrease in shocks, hence the emission detected in the L1157-B1 shock is a probe of pre-shock chemistry. Molecular ions such as \\hocop\\ and \\sop, instead, are believed to be effective tracers of shocked gas, as their abundance is predicted to be very low in the quiescent gas ($\\le 10^{-10}$) but could be strongly enhanced in shocks following the release of \\coo\\ and S-bearing molecules from dust grains mantles \\citep[e.g., ][]{minh91,turner92,turner94,deguchi06}. However, there are no observations of \\hocop\\ and \\sop\\ undoubtedly associated with shocks.    In this paper, for the first time a complete census of molecular ions in a shock is given. This is obtained by means of a highly-sensitive unbiased spectral survey of the prototypical  protostellar bow-shock L1157-B1  in the millimeter and submillimeter range, executed with the IRAM-30m telescope as part of the ASAI Large Program\\footnote{http://www.oan.es/asai} (Astrochemical Surveys At Iram, Lefloch, Bachiller et al. in preparation) and with {\\it Herschel}/HIFI as part of the Guaranteed Time Key Project CHESS\\footnote{http://chess.obs.ujf-grenoble.fr} (Chemical HErschel Surveys of Star forming regions, \\citealt{ceccarelli10}). The observations are presented in Sect.~\\ref{sect:observations}, while Sect.~\\ref{sect:l1157} summarizes the main properties of the L1157 outflow and the B1 bow-shock. In Sect.~\\ref{sect:results} we analyze the line profiles of the detected molecular ions and we perform a Large Velocity Gradient (LVG) analysis to infer the physical conditions of the emitting gas, and the abundance of molecular ions. Finally, in Sect.~\\ref{sect:astrochem} we compare the observed abundances with steady-state abundances in the cloud and their evolution in the shock computed using the chemical code Astrochem. This allows us to distinguish between molecular ions which are tracing pre-shock chemistry, and those which are a useful probe of the shock properties and chemistry. Our conclusions are summarized in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}   In this paper we presented a comprehensive study of the chemistry of molecular ions in protostellar shocks. As part of the ASAI and CHESS projects, we performed a census of molecular ions in the chemically rich protostellar shock L1157-B1 by means of an unbiased and high sensitivity ($\\sim 1.5$~mK in the 3~mm band) spectral survey executed with the IRAM-30m telescope and {\\it Herschel}/HIFI. % From our analysis we obtained the following results: \\\\  \\begin{itemize}  \\item[-] we report emission from \\hcop, \\nnhp, \\hcsp\\ and, for the first time in a protostellar shock, from \\hocop\\ and \\sop.\\\\  \\item[-] all the detected lines peak at blueshifted velocity, $\\sim 0.5-3$ \\kms\\ with respect to systemic (V$_{\\rm sys} = + 2.6$~\\kms), and have a line width of $\\sim 3-7$~\\kms. A higher velocity component (V up to $- 40$~\\kms) associated with the primary jet is detected in the \\hcop~1--0 line.\\\\  \\item[-] the lines profiles can be decomposed into different velocity components following an exponential intensity-velocity relation as found for the CO lines \\citep{lefloch12}. The bulk of the emission is associated with the extended outflow cavities and originates from a gas with temperature T$_{\\rm kin} \\sim 20-70$ K, and density \\denshh $\\sim 10^5$ \\cmc, in agreement with what found for the CO lines.\\\\  \\item[-] inferred \\hcop\\ and \\nnhp\\ abundances (\\xhcop~$\\sim0.7-3~10^{-8}$, \\xnnhp~$\\sim0.4-8~10^{-9}$) are in agreement with on-source values \\citep{bachiller97}, suggesting that these species are not enhanced in the shock. This is further confirmed by the fact that the observed abundances are in agreement with steady-state abundances in the cloud (i.e. at T$_{\\rm kin} = 10$~K, \\denshh~$= 10^4$ \\cmc). The decrease of their abundance due to gas compression in the shock is compensated for by the raise in temperature, and predicted values match the observations for high cosmic rays ionization rate $\\zeta = 3 \\ 10^{-16}$ s$^{-1}$. Hence, \\hcop\\ and \\nnhp\\ are a fossil-record of the gas chemistry before the arrival of the shock.\\\\  \\item[-] \\hocop, \\sop, and \\hcsp\\ abundances  (\\xhocop~$\\sim10^{-9}$, \\xsop~$\\sim8~10^{-10}$, \\xhcsp~$\\sim3-7~10^{-10}$), instead, are 1--2 orders of magnitude larger than predicted steady-state abundances in the cloud. On the other hand, they are strongly enhanced in the shock if \\coo, S, H$_2$S, and OCS abundances are increased with respect to steady-state due to release off the dust grains in the shock. Hence, as suggested by previous studies \\citep[e.g. ][]{minh91,deguchi06,turner92,turner94}, these species are effective shock tracers.\\\\  \\item[-] the observed abundances allow constraining the amount of \\coo\\ and sulphur-bearing species released in the shock. To reproduce \\hocop\\ observations, \\coo\\ should be enhanced up to abundances $\\ge 2 \\ 10^{-5}$, i.e. two order of magnitude larger than its gas-phase abundance in the quiescent gas. This is in agreement with previous estimate of the \\coo\\ content on dust mantles. To reproduce \\sop, instead, the sulphur released in the form of H$_2$S molecules should be $\\ge 2 \\ 10^{-7}$. This value is in agreement with the H$_2$S abundance in B1 estimated by \\citet{bachiller97}.\\\\  \\item[-] as noted in previous studies, CS and \\hcsp\\ abundances in the quiescent gas in dark clouds are very low.  However, the observed \\hcsp\\ and CS abundances (\\xhcsp~$\\sim 3-7~10^{-10}$, \\xcs~$\\sim 2~10^{-7}$) are matched if  OCS is assumed to be one of the main sulphur career on dust grains, as suggested by \\citet{wakelam04,wakelam05} and \\citet{codella05}. Hence its abundance is enhanced to values $\\ge 2 \\ 10^{-6}$ in the shock. Such high \\xocs\\ values, however, are at least one order of magnitude larger than estimated by \\citet{bachiller97} and are lowered to the observed values only on timescales $\\ge 10^4$ years, i.e. much larger than the shock age (t$_{\\rm shock} \\sim 2000$ years). This difficulty in simultaneously reproducing CS, \\hcsp, and OCS abundances was already reported by other authors and indicates that the chemistry of sulphur-bearing species is not yet fully understood \\citep{doty04,wakelam04,wakelam05}. In this regard, laboratory experiments are required to verify the rates of ion-molecule reaction producing \\hcsp, i.e. (\\hcop, \\hhhp, H$_3$O$^{+}$) + CS. \\\\   \\end{itemize}"
    },
    "1402/1402.4739_arXiv.txt": {
        "abstract": "{The attenuation of atmospheric Cherenkov photons is dominated by two processes: Rayleigh scattering from the molecular component and Mie scattering from the aerosol component. Aerosols are expected to contribute up to 30\\,Wm$^{-2}$ to the emission profile of the atmosphere, equivalent to a difference of $\\sim 20^\\circ$C to the clear sky brightness temperature under normal conditions. Here we investigate the aerosol contribution of the measured sky brightness temperature at the H.E.S.S. site; compare it to effective changes in the telescope trigger rates; and discuss how it can be used to provide an assessment of sky clarity that is unambiguously free of telescope systematics.} ",
        "introduction": "The atmosphere is the most important part of the detector in ground-based gamma-ray astronomy, but it is also the part that has the greatest systematic uncertainty and over which we have the least control. It falls upon us to instead monitor and characterise the atmospheric conditions at the time of observations so that we can either feed this information into Monte Carlo simulations or reject data when conditions go out of acceptable parameters.  Cherenkov light is generated in the upper atmosphere and will only reach a ground based Cherenkov telescope if not attenuated through the process of Rayleigh scattering on the molecular component of the atmosphere, or Mie scattering on the aerosol component (variously dust, silicates, pollens, smoke, etc). The molecular component of the atmosphere tends to change relatively slowly, through seasonal variations; whereas the aerosol component can change much more rapidly, depending on wind conditions for example. It becomes vitally important to characterise this aerosol component of the atmosphere through regular monitoring taken concurrently with telescope observations. A lidar is generally used to measure the atmospheric transmission (eg \\cite{bib:nolan}) from backscattered laser light. At the H.E.S.S. site a lidar centred at 355 and 532\\,nm has been running in conjunction with observations since mid-2011 \\cite{bib:vasileiadis}. Whilst lidars are excellent instruments for determining the presence of aerosols they are not without complications. Firstly a lidar, due to geometric viewing considerations, only becomes effective above a minimum altitude, whereas the aerosol concentration will always be greatest close to the ground. Secondly, in order to obtain a transmission profile relevant to the Cherenkov spectrum the laser wavelengths are close to the peak in the emission, this means the H.E.S.S. lidar is operated only between observing runs to avoid any light contamination to the telescope images. In this paper we look at utilising another piece of the H.E.S.S. atmospheric monitoring equipment, the all sky scanning infrared radiometer, to fill in some of this missing information.  The radiometer measures an equivalent brightness temperature for the atmosphere and correlates this with the atmospheric clarity conditions. The atmosphere is split into regions according to its temperature behaviour. The troposphere is the lowest, most dense, part of the atmosphere -- where most of the weather happens -- and is characterised by a linear decline in temperature with increasing altitude and vertical mixing. The molecular density profile falls off exponentially, with a scale height of a $\\sim$few km; the vertical air motion in this region also mixes in the larger aerosols which have a smaller scale height of order a $\\sim$km. The molecular component is an inefficient black-body radiator in the 8-14$\\mu$m region of the spectrum, water vapour and aerosols are slightly more efficient and clouds are very efficient. This makes an infra-red radiometer an effective cloud monitor (and this is the primary function of this instrument), with clouds showing up as a large brightness temperature compared to a relatively ``cold\" sky \\cite{bib:buckley}. H.E.S.S. employ Heitronics KT19.82 radiometers with a Germanium lens of 2$^\\circ$ field of view to monitor for the presence of clouds, with each telescope having a paraxially mounted unit and a further independent unit continuosly scanning the whole sky.  The infra-red luminosity of the sky ($L_{\\mathrm{sky}}$) is actually a collective sum of the emission of a number of different constituent parts \\begin{eqnarray*} L_{\\mathrm{sky}} = \\epsilon_{l}\\sigma T_{\\mathrm{lens}}^4 + \\epsilon_{wv}\\sigma \\int T_{wv}^4 + \\epsilon_{a}\\sigma \\int T_{a}^4 + \\epsilon_{m}\\sigma \\int T_{m}^4 + \\ldots \\end{eqnarray*} where $\\epsilon$ denotes the emissivity of the lens ($l$), the water vapour ($wv$), the aerosols $a$, and the molecular ($m$) profiles of the atmosphere, respectively and $\\int T$ is the relevant integrated temperature profile in the line of sight. As the lower altitudes are also the warmest these will contribute the most to the observed infrared radiation budget, this is also the region that will be most affected by the varying aerosol concentration. According to \\cite{bib:dalrymple} the aerosol component can contribute up to 30\\,Wm$^{-2}$ to the bolometric luminosity of a clear sky\\footnote{which can mean the difference between a brightness temperature of -56$^\\circ$C or -70$^\\circ$C in the presence or absence of aerosols respectively}. This leads to the prospect of changing aerosol conditions leading to a noticeable change in the sky brightness temperature ($L_{\\mathrm{sky}} = \\sigma T^{4}_{\\mathrm{sky}}$) measurements.  \\begin{figure*}[!ht] \\centering \\includegraphics[width=0.4\\textwidth]{20110825-T.png} \\includegraphics[width=0.4\\textwidth]{20110825-L.png} \\includegraphics[width=0.4\\textwidth]{20110829-T.png} \\includegraphics[width=0.4\\textwidth]{20110829-L.png} \\caption{Scanning radiometer sky brightness temperature measurements as a function of zenith angle. The top plots are for a clear night with no discernable aerosol boundary layer; the bottom for a night with a boundary layer up to $\\sim3.5$\\,km. The left plots show the all sky temperature values as a function of zenith and azimuth angle, the right plots are histograms of the accumulated luminosity values as a function of zenith angle. } \\label{fig:T} \\end{figure*}  \\begin{figure*}[!t] \\centering \\includegraphics[width=\\textwidth]{20110824_VsZenith.png} \\caption{The telescope trigger rate and sky brightness temperature as a function of zenith angle for a ``clear sky''.} \\label{fig:clear} \\end{figure*} ",
        "conclusions": "The correlation between sky brightness temperature and lidar determined aerosol concentration has been shown before \\cite{bib:daniel} on a limited data set. Here the study has been expanded to larger, year-long, dataset and once again a positive correlation of sky brightness temperature and aerosol concentration is seen. This also corresponds to a clear inverse correlation with Cherenkov telescope trigger rates. This shows that a radiometer can be a useful instrument not just as a cloud monitor, but also as a sensitive measure of atmospheric clarity. As the radiometer will be most sensitive to the warmest aerosols it is a useful instrument in the low altitude region where a lidar is not sensitive due to it not having reached full geometric overlap. An infrared radiometer is also a completely passive instrument, allowing it to take fully contemporaneous data with a Cherenkov telescope without affecting observations. It can also take data during the daytime, allowing a forecast of that night's observing conditions to be easily made.  \\vspace*{0.5cm} \\footnotesize{{\\bf Acknowledgment: }{We would like to acknowledge the support of the technical staff at Durham and the H.E.S.S. telescopes in the operation of the radiometer and the people at LUPM in the running of the lidar.}}"
    },
    "1402/1402.4814_arXiv.txt": {
        "abstract": "We use SDSS+\\textit{GALEX}+Galaxy Zoo data to study the quenching of star formation in low-redshift galaxies. We show that the green valley between the blue cloud of star-forming galaxies and the red sequence of quiescent galaxies in the colour-mass diagram is not a single transitional state through which most blue galaxies evolve into red galaxies. Rather, an analysis that takes morphology into account makes clear that only a small population of blue early-type galaxies move rapidly across the green valley after the morphologies are transformed from disk to spheroid and star formation is quenched rapidly. In contrast, the majority of blue star-forming galaxies have significant disks, and they retain their late-type morphologies as their star formation rates decline very slowly. We summarize a range of observations that lead to these conclusions, including UV-optical colours and halo masses, which both show a striking dependence on morphological type. We interpret these results in terms of the evolution of cosmic gas supply and gas reservoirs. We conclude that late-type galaxies are consistent with a scenario where the cosmic supply of gas is shut off, perhaps at a critical halo mass, followed by a slow exhaustion of the remaining gas over several Gyr, driven by secular and/or environmental processes. In contrast, early-type galaxies require a scenario where the gas supply and gas reservoir are destroyed virtually instantaneously, with rapid quenching accompanied by a morphological transformation from disk to spheroid. This gas reservoir destruction could be the consequence of a major merger, which in most cases transforms galaxies from disk to elliptical morphology, and mergers could play a role in inducing black hole accretion and possibly AGN feedback. ",
        "introduction": "\\label{sec:intro}  \\begin{keywords} galaxies: evolution; galaxies: active; galaxies: spiral; galaxies: elliptical and lenticular, cD \\end{keywords}  Ever since the discovery of the bimodality in galaxy colour in the galaxy colour-magnitude and colour-mass diagrams from large-scale surveys \\citep{2001AJ....122.1861S, 2004ApJ...600..681B, 2006MNRAS.373..469B}, the colour space between the two main populations --- the \\textit{green valley} --- has been viewed as the crossroads of galaxy evolution. The galaxies in the green valley were thought to represent the transition population between the blue cloud of star-forming galaxies and the red sequence of quenched, passively evolving galaxies \\citep[e.g.,][]{2004ApJ...608..752B, 2007ApJS..173..293W, 2007ApJS..173..315S, 2007ApJS..173..342M, 2007ApJ...665..265F,2011ApJ...736..110M,2012ApJ...759...67G}. Roughly speaking, all galaxies were presumed to follow similar evolutionary tracks across the green valley, with a fairly rapid transition implied by the relative scarcity of galaxies in the green valley compared to the blue cloud or red sequence.  The intermediate galaxy colours of green valley galaxies have been interpreted as evidence for the recent quenching of star formation \\citep{2007ApJS..173..267S}. The clustering of active galactic nuclei (AGN) host galaxies in the green valley further suggested a role for AGN feedback in particular \\citep[e.g.,][]{2007ApJ...660L..11N,2008A&A...490..905H, 2008ApJ...675.1025S}. Galaxies in the green valley have specific star formation rates (sSFR) lower than the ``main sequence\" of star formation in galaxies, which is a tight correlation between galaxy stellar mass and star formation rate, presumably as a result of quenching \\citep[e.g.][]{2004MNRAS.351.1151B, 2007A&A...468...33E, 2007ApJS..173..267S, 2007ApJ...660L..43N, 2010ApJ...721..193P, 2011A&A...533A.119E,2012ApJ...752...66L, 2012ApJ...745..149L}). Most star-forming galaxies live on the main sequence, so tracing the populations leaving the main sequence -- those with lower sSFRs -- probes the quenching mechanism(s) and, as \\cite{2010ApJ...721..193P} showed, there may be at least two very different quenching processes.  Ultraviolet light comes predominantly from newly formed massive stars, which makes UV observations an excellent probe of the current rate of star formation. In this paper, we use UV-optical colours from \\textit{GALEX} photometry \\citep[see e.g.,][]{2007ApJS..173..342M} to investigate the rate at which galaxies are decreasing their sSFR (i.e., how rapidly they change colour), and whether this correlates with morphology.  Essentially, we use galaxy colours as stellar population clocks, an approach first conceived and applied by Tinsley and collaborators \\citep[][]{1968ApJ...151..547T,1976ApJ...203...52T,1978ApJ...221..554T}.  We interpret the evolutionary tracks of disks in terms of the gas supply and how star formation depletes the gas reservoir \\citep{1959ApJ...129..243S}. Interestingly, what had appeared to be outliers from the general parent galaxy population --- namely, blue early types \\citep{2009MNRAS.396..818S} and red late types \\citep{2010MNRAS.405..783M} --- far from being curiosities, are instead a valuable clue to galaxy evolution.  Morphology has not previously been a major ingredient in interpretations of the colour-mass diagram. Now that reliable morphological classifications have been made possible by citizen scientists in the Galaxy Zoo project \\citep{2008MNRAS.389.1179L,2011MNRAS.410..166L},  we are able to investigate the relation of galaxy morphology to colour and mass. We also consider galaxy content and environment, now that standardized information is available for large galaxy samples \\citep[e.g.][]{2004ApJ...600..681B,2007ApJ...671..153Y,2010ApJ...721..193P}.  The Galaxy Zoo data have already enabled many insights about the link between galaxy evolution and colour \\citep{2009MNRAS.396..818S, 2010MNRAS.405..783M} and about the link between galaxy evolution and environment \\citep[e.g.,][]{2009MNRAS.393.1324B,2009MNRAS.399..966S}, mergers \\citep{2010MNRAS.401.1552D,2010MNRAS.401.1043D, 2012MNRAS.419...70K,2012ApJ...753..165T}, unusual galaxy types \\citep{2009MNRAS.399.1191C, 2009MNRAS.399..129L, 2012MNRAS.420..878K}, and specific morphological features such as bars \\citep[e.g.,][]{2011MNRAS.411.2026M, 2011MNRAS.415.3627H, 2012MNRAS.423.1485S, 2012MNRAS.424.2180M, 2013ApJ...779..162C, 2014MNRAS.tmp...97M}. In this paper, Galaxy Zoo morphologies provide the key to understanding that early- and late-type galaxies,  even those with similar green optical colours, follow distinct evolutionary trajectories involving fundamentally different quenching modes.  Throughout this paper, we use a standard $\\Lambda$CDM Cosmology ($\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $\\mathrm{H_{0}} = 70~\\mathrm{km^{-1}~s^{-1}}$), consistent with observational measurements \\citep{2011ApJS..192...18K}. All magnitudes are in the AB system.    \\begin{figure*} \\begin{center}  \\includegraphics[width=0.99\\textwidth]{gallery.pdf}  \\caption{Example $gri$ SDSS images (51.2\\arcsec\\ $\\times$ 51.2\\arcsec), ordered by Galaxy Zoo classification. In the \\textit{top}, \\textit{middle} and \\textit{lower} rows are early-, indeterminate- and late-type galaxies, respectively. In each row, we show (from left to right) three blue-cloud, three green-valley and three red-sequence galaxies. The indeterminate-type galaxies are mostly composite bulge-disk systems that more closely resemble the late-type galaxies than the purely spheroidal early types. For this reason, it is not surprising that they mostly follow the late-type galaxies in their quenching behaviour.}  \\label{fig:gallery}  \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center}  \\includegraphics[width=0.99\\textwidth]{colour_mass_gv.pdf}  \\caption{The $u-r$ colour-mass diagram for our sample. In the {\\it top left}, we show all galaxies, whereas on the {\\it right}, we show the early-type ({\\it top}) and late-type galaxies ({\\it bottom}); green lines show the green valley defined by the all-galaxy diagram. This figure illustrates two important findings: (1) Both early- and late-type galaxies span almost the entire $u-r$ colour range. Visible in the morphology-sorted plots are small numbers of blue early-type ({\\it top}) and red late-type ({\\it bottom}) galaxies \\citep[e.g.,][]{2009MNRAS.396..818S,2010MNRAS.405..783M}. (2) The green valley is a well-defined location only in the all-galaxies panel ({\\it upper left}). Most early-type galaxies occupy the red sequence, with a long tail (10\\% by number) to the blue cloud at relatively low masses; this could represent blue galaxies transiting rapidly through the green valley to the red sequence. More strikingly, the late types form a single, unimodal distribution peaking in the blue (these are the main sequence star-formers) and reaching all the way to the red sequence, at higher masses, with no sign of a green valley (in the sense of a colour bimodality). The contours on this Figure are linear and scaled to the highest value in each panel. }  \\label{fig:colour_mass_gv}  \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  We have used a series of observational results, informed by morphological classifications, to develop a broad picture of how and why early- and late-type galaxies at $z\\sim0$ transform from blue star-forming galaxies into passively evolving red galaxies. Here we review the evidence presented in Section \\ref{sec:gv}, then discuss our interpretation.  We first verified that dust extinction, while present, is not the main reason most galaxies have green colours. We then assessed the recent star formation properties of green valley galaxies, as traced by emission lines and by the location of green valley galaxies compared to the main sequence. Both indicators show that low star formation rates are the reason green valley galaxies --- regardless of morphology --- exhibit green optical colours.  We then considered the UV-optical colours of galaxies in the optical ($u-r$) green valley, the UV (from \\textit{GALEX}) being more sensitive to the youngest stars than purely optical colours. We found that, when early-type galaxies are in the (optical) green valley, their UV-optical colours are much redder than those of late-type green valley galaxies. This supports the idea that the two morphological classes have fundamentally different recent star formation histories (\\textit{i.e.} different UV-u colors), even after the initial stellar populations have aged in a similar way (\\textit{i.e.} optical colors are green). We used model evolutionary tracks to show that this difference corresponds to very different time scales on which star formation declines: in early-type galaxies, the quenching time scale is much shorter ($<250$~Myr) than in late-type galaxies ($>1$~Gyr).  All this evidence leads to a coherent physical picture where a quenching event destroys the (quasi) equilibrium state of main sequence star-forming galaxies, moving them off the main sequence into the green valley. The data clearly show that {\\it the time scale for this colour evolution is strongly tied to morphology}. We now argue that two quenching pathways --- for early-type and for late-type galaxies --- can be tied to two different scenarios for destroying gas reservoirs, i.e., to two different physical quenching mechanisms. Because most star-forming galaxies start on the main sequence of star formation, where a simple regulator balances cosmological inflows of gas with outflows \\citep{2013arXiv1303.5059L}, we take the main sequence as the starting point of the evolutionary models discussed below.  \\subsection{The evolution of late-type galaxies}  We have shown that the star formation histories of late-type galaxies are consistent with a very gradual quenching, corresponding to an exponential time scale of a gigayear or more. A natural explanation is that the baryon supply is disrupted, so that the star formation rate declines slowly in response to the depletion of the gas reservoir. The disruption could be due to halo mass quenching \\citep[e.g.,][]{2006MNRAS.370.1651C}, the end of cold streams and the development of a hot halo due to shocks \\citep[]{2006MNRAS.368....2D, 2009Natur.457..451D}, or simply the continued expansion of the virial radius to a point where cooling from the halo to the disk becomes inefficient (i.e., too slow). While other environmental processes may also play a role --- for example, ram pressure, stripping, or harassment (see \\citealt{2011nha..confE...8T} and \\citealt{2013pss6.book..207V} for recent reviews) --- these local considerations cannot explain the global behavior of late-type galaxies as a class.  The rate at  which star formation exhausts gas can be written in a dynamical form as \\citep[][]{1998MNRAS.295..877G, 2003MNRAS.343...75H, 2011MNRAS.415.3798K}:  \\begin{equation} \\psi = \\frac{\\epsilon M_{\\rm gas}}{\\tau_{\\rm dyn}} , \\label{eq:schmidt} \\end{equation}  \\noindent where $\\psi$ is the SFR, $M_{\\rm gas}$ is the gas reservoir available for star formation, $\\tau_{\\rm dyn}$ is the dynamical time scale, and $\\epsilon$ is the (fixed) efficiency of star formation \\citep[we assume the universal value $\\epsilon=2$\\%; e.g.,][]{1998ApJ...498..541K}. The dynamical time scale for the disk of a massive galaxy is on the order of $\\tau_{\\rm dyn} \\sim 50-300$ Myr. The mass of the gas reservoir is a free parameter set by the initial conditions. Once the initial gas reservoir is specified and a reasonable dynamical time scale is chosen, Eqn.~\\ref{eq:schmidt} can be solved to describe the evolution of the system. The e-folding time for the declining star formation rate is $\\tau_{\\rm dyn}/ \\epsilon$, which is several Gyr for the adopted parameters\\footnote{This time scale for gas depletion corresponds to what happens in the ``bathtub'' model of \\cite{2013arXiv1303.5059L} when the gas supply to the reservoir is shut off.}. That is, galaxies with a gas reservoir that does not get replenished from outside will continue to form stars for a very long time, peeling off from the main sequence of star formation very slowly, as previously argued by \\cite{2011MNRAS.415.3798K}. In Figure \\ref{fig:tracks}, we show how the specific star formation rates, $u-r$ colours, and gas fractions evolve according to Eqn.~3 for a range of quenching time scales.  The evolution of late-type galaxies corresponds to the longer time scale tracks in Figure \\ref{fig:tracks}. This means there is enough time to develop a significant population of late types in the green valley and, eventually, the red sequence. Late-type galaxies in the green valley should still have substantial gas reservoirs that fuel ongoing star formation. Initial data from HI data support this (Section \\ref{sec:HI_gas}).  The start of quenching may occur at (very) high redshift, as it can take several Gyrs for the effects of a slowly declining star formation rate to become apparent. In this context, red spirals \\citep[e.g.,][]{1976ApJ...206..883V, 2009MNRAS.393.1302W, 2010MNRAS.405..783M} are simply late-type galaxies whose cutoff happened relatively early. The observation of \\cite{2010MNRAS.405..783M} and \\cite{2012A&A...543A.132C} that red spirals still show signs of low-level star formation make sense --- they are simply very far along the exponential decline in SFR off the main sequence.  The reddest colours in late-type galaxies preferentially occur in high mass haloes (Fig.~\\ref{fig:gv_env}), supporting the suggestion that the gas supply is somehow regulated by environmental factors.  Of course, other factors besides halo mass can increase the pace of gas depletion following the original quenching event --- and this must happen: Figure \\ref{fig:evol1} shows some objects with quenching time scales slightly shorter than that expected from canonical parameters.  For example, bars naturally shorten gas consumption time scales, by driving gas into the central regions of galaxies where it can be consumed more quickly. \\cite{2011MNRAS.411.2026M, 2012MNRAS.424.2180M} found that the incidence of bars in massive late-type galaxies increases as they become redder. This could be a side effect of quenching --- models show that bars form more quickly in gas-poor disks \\citep[e.g.,][]{2013MNRAS.429.1949A} --- but if secular evolution along a bar is an efficient process, the presence of a bar could also speed up the quenching time such that the transition happens more rapidly than the several Gyrs expected from the simple Schmidt law. \\cite{2012MNRAS.424.2180M} show that high bar fractions in spirals are associated with lower HI gas content, and there are hints that at a fixed HI content barred galaxies are optically redder than unbarred galaxies. Several works (\\citealt{2008ApJ...675.1141S, 2010MNRAS.409..346C}; Melvin et al. in prep) show that the fraction of disk galaxies with bars also seems to have been increasing since $z\\sim1$ \\citet{2008ApJ...675.1141S, 2010MNRAS.409..346C}. Thus, secular processes like bar formation may be intimately tied to quenching in late-type galaxies; even if they appear after the initial quenching event, they would help drive out the remaining gas reservoir, thus accelerating the rate at which late types become fully gas- and star-formation free.  Environment as a late-stage accelerator or modifier of quenching is naturally supported by the strong effect of halo mass on the presence of green and red late types (Section \\ref{sec:env_gv}). Other recent observational studies have argued that environment or satellite quenching should act fast, at least once it begins \\citep[e.g.,][]{2012ApJ...746..188M, 2013MNRAS.432..336W}.  Other environmental effects can accelerate the quenching process by further removing gas but not too quickly, as we do not see late-type galaxies with short quenching timescales. We note that these modulating effects (secular processes and environment) mean that most likely we cannot accurately reconstruct the time of quenching for present-day green and red (i.e., off-main sequence) late types.  Now we consider the possible role of AGN feedback in the evolution of late-type galaxies.The majority of black holes growing locally and at high redshift are hosted in late-type galaxies, clearly driven by secular processes rather than major mergers \\citep[e.g.,][]{2010ApJ...711..284S,2011ApJ...727L..31S, 2012MNRAS.425L..61S, 2011ApJ...726...57C, 2012ApJ...744..148K, 2012ApJ...761...75S, 2012ApJ...758L..39T}. The fraction of green valley galaxies that host AGN is higher than in blue star-forming galaxies (Fig.~11). This could be due in part to an observational bias (the difficulty of seeing AGN signatures in SDSS spectra against a more luminous stellar component), but sensitive multiwavelength searches for  ``missed'' AGN in nearby star-forming galaxies \\citep[e.g.,][]{2010MNRAS.406..597G} do not find enough to balance the numbers. Moreover, hard X-ray selection should not be confused by star formation, yet this technique also detects few to no (high luminosity) AGN in the blue cloud \\citep{2009ApJ...692L..19S}. We conclude the prevalence of AGN in green valley galaxies is likely real rather than a selection effect.  The alternative is that the onset of quenching in late-type galaxies leads to increased black hole growth about 1~Gyr later.  The delay could be explained naturally by the time required for accreting matter to lose angular momentum. In this case, black hole feeding would be due to secular processes (since major mergers would disrupt the disk morphology) and could not significantly speed up the quenching (or the AGN hosts would transition much more rapidly across $NUVur$ colour space); it would be a consequence of the host galaxy quenching, not the cause. Perhaps a declining sSFR favors black hole growth: \\cite{2007ApJ...671.1388D} suggested that the absence of strong stellar feedback (core collapse supernovae, O-star winds) makes it easier to transport gas to the center, and also that mass loss from Asymptotic Giant Branch stars could easily fuel black hole accretion because of both the larger amount of material and its lower velocity.  \\cite{2010ApJ...721..193P} identified two distinct quenching mechanisms at work in galaxies: environment and mass quenching (without identifying the physical mechanisms). From the data presented in Section \\ref{sec:env_gv}, one might speculate that early-type galaxies are mass quenched (no strong dependence on environment) and late types are environment quenched (strong dependence on environment). A test of this hypothesis would be the mass functions of green valley early- and late-type galaxies: if the quenching pathway we identify in early types is mass quenching (see below), then the green valley early types should have the same mass function as the passive galaxies (in terms of $M*$ and faint end slope $\\alpha_{\\rm s}$, the normalisation $\\Phi*$ will depend on transition time scale); similarly, if the quenching pathway we identify in late types is environment quenching, then the green valley late types should have the same mass function as the star-forming galaxies.   \\subsection{The evolution of early-type galaxies}  The end of star formation in early-type galaxies in the local Universe proceeds in a fundamentally different fashion than in the late-type galaxies. In local early types, quenching occurs in low-mass galaxies and is marked by a very rapid shutdown of star formation on time scales consistent with instantaneous suppression, or at most $\\tau_{\\rm quench} \\sim 250$ Myr \\citep[c.f.][]{2007MNRAS.382.1415S}. This rapid suppression is inconsistent with the simple gas exhaustion scenario outlined for late-type galaxies; the Schmidt law does not allow star formation to deplete a substantial gas reservoir so rapidly \\citep[a point made strongly by][]{2011MNRAS.415.3798K}. The quenching mechanism should also be linked to the destruction of the disk morphology of the likely progenitor(s).  A major merger \\citep{2010ApJ...714L.108S} could simultaneously transform the galaxy morphology from disk to spheroid and cause rapid depletion of the cold gas reservoir. Deep imaging of blue early types does reveal tidal features indicative of a major merger \\citep{2010ApJ...714L.108S} but \\cite{2012MNRAS.420.1684W} have shown that galaxies with post-starburst spectral features already have early-type morphologies, emphasizing that the morphological transformation must be rapid.  Studies of gas kinematics (molecular, ionised) show that most early-type galaxies, and specifically those in the green valley, have gas with an external origin, most likely due to merger activity \\citep{2006MNRAS.366.1151S,2010MNRAS.402.2140S,2013MNRAS.429..534D,2003ApJ...597L.117K}. These blue precursors must have similar masses but bluer (intrinsic) colours than the AGN early-type hosts (see Figure \\ref{fig:colour_mass_gv_nodust} and \\citealt{2007MNRAS.382.1415S,2009MNRAS.396..818S}).  Observations of a small sample of early types along this evolutionary path from the blue cloud to the red sequence shows that the large molecular gas reservoirs of blue early-type galaxies do disappear rapidly,  at a rate significantly faster than can be explained by star formation alone \\citep{2009ApJ...690.1672S}.  We note that, as with the green valley late types, the quenching event --- i.e., the point at which the gas fuelling star formation was destroyed --- had to occur \\textit{before} the galaxy reaches the green valley. Given the short time between quenching and arrival in the green valley (only a few hundred Megayears), it may be easier to identify this progenitor population than in late types. So far, \\cite{2009ApJ...690.1672S} found that green valley (Seyfert AGN) early-type galaxies were undetected in molecular gas and that the gas disappeared rapidly during the transition from purely star-forming to an AGN+star formation mixed phase.  Naturally, today's galaxies host both stellar populations formed {\\it in situ} and those brought in via progenitors in minor and major mergers. In this general picture, star formation in early-type galaxies ceases and then does not re-start except for a minor frosting of new stellar populations \\citep[seen mostly in the ultraviolet; e.g.,][]{2005ApJ...619L.111Y, 2007ApJS..173..512S, 2007ApJS..173..619K}, except when a new disk forms through the acquisition of a sufficient supply of cold gas, at which point the galaxy would rejoin the main sequence.  What about the role of AGN feedback? Rapid depletion of a large fraction of the available gas could be caused by a vigorous starburst, as star formation uses up cold gas and stellar processes create strong winds. Simulations show that merger-induced starbursts can lead to enhanced star formation as disk destruction leads to the inflow of gas to the (new) galaxy centre, all on relatively short dynamical time scales \\citep[e.g.,][]{1996ApJ...471..115B}.  However, in this starburst-induced scenario, even a short depletion time scale will still yield a remnant system that is not entirely quenched. The Schmidt Law forces the star formation back to an exponentially decaying state, albeit with a shorter time scale, which in turn means that the galaxy retains gas, and therefore continues forming stars.  Only by adding AGN feedback, which can destroy the gas reservoir, do simulations show a genuine total quenching of star formation \\citep[e.g.,][]{2005ApJ...620L..79S}.  Is it possible that AGN feedback alone destroys the gas reservoir during merger? It is true that photoionization happens almost instantaneously, so a luminous AGN could destroy the cold gas reservoir almost instantaneously. However, we would then expect detectable AGN radiation before the host galaxies reach the green valley, while they are still in the blue, star-forming phase, contrary to what is observed \\citep[][]{2009ApJ...692L..19S, 2010MNRAS.406..597G}. If AGN feedback were responsible for the rapid gas reservoir destruction, then during this phase it must be either very short-lived, heavily obscured or radiatively inefficient \\citep{2009ApJ...690.1672S,2012arXiv1206.2661S}.  A radiatively inefficient accretion flow could drive a kinetic outflow (jet, disk wind, or other outflow), analogous to what X-ray binaries do \\citep[e.g.,][]{2003MNRAS.345L..19M,2006MNRAS.372.1366K, 2006Natur.444..730M, 2010Natur.466..209P}. In this early phase, the AGN would be outshone by the declining starburst; the place to look for evidence of this kind of kinetic feedback would be at the transition from star forming to composite spectrum in the blue early-type galaxy population.  We note that dwarf ellipticals are not in our sample and their quenching pathways may be very different \\citep[e.g.][]{2008ApJ...674..742B}    \\begin{figure*} \\begin{center}  \\includegraphics[width=0.99\\textwidth]{late_type_schema_v7.pdf}   \\caption{Diagram outlining the scenario presented here for star-formation quenching in late-type galaxies. Once the external supply of gas to the galaxy's reservoir is cut off, the galaxy will leave the main sequence, and the SFR will decay exponentially, with a long quenching time scale $\\tau_{\\rm quench}$. While no longer on the main sequence, the galaxy nevertheless continues to convert gas to stars and thus to increase its stellar mass. Absent a major merger, this evolutionary pathway eventually  produces a passive, red, late-type galaxy. We see radiation from black hole growth during stage (3), long after the quenching event actually took place. }  \\label{fig:lt_cartoon}  \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center}  \\includegraphics[width=0.99\\textwidth]{early_type_schema_v6.pdf}   \\caption{Diagram outlining the scenario presented here for star-formation quenching in early-type galaxies. A dramatic event, likely a gas-rich merger of main sequence star-forming galaxies, causes the near-instantaneous destruction of the (newly formed) early-type galaxy gas reservoir, via stellar processes and/or AGN feedback. As long as the gas reservoir is not resupplied or the AGN keeps the gas hot, star formation ceases and no further stellar mass is formed. An AGN phase is visible for 100s of Myr as the stars in the early-type host galaxy age rapidly through the green valley. }  \\label{fig:et_cartoon}  \\end{center} \\end{figure*}"
    },
    "1402/1402.4449_arXiv.txt": {
        "abstract": "We use photometry from the recent AllWISE Data Release of the {\\it Wide-field Infrared Survey Explorer} ({\\it \\hspace{-1.5pt}WISE}\\,) of 129 calibration stars, combined with prior distances obtained from the established $M_V$--[Fe/H] relation and {\\it Hubble Space Telescope} trigonometric parallax, to derive mid-infrared period--luminosity relations for RR Lyrae pulsating variable stars. We derive relations in the $W1$, $W2$, and $W3$ wavebands (3.4, 4.6, and 12 $\\mu$m, respectively), and for each of the two main RR Lyrae sub-types (RRab and RRc). We report an error on the period--luminosity relation slope for RRab stars of 0.2. We also fit posterior distances for the calibration catalog and find a median fractional distance error of 0.8 per cent. ",
        "introduction": "As calibratable standard candles, RR Lyrae pulsating variable stars provide a means to obtain highly precise (1-2 per cent error) distance measurements within the range of 1-100 kpc. These old (age $\\gtrsim10\\times10^9$ yr) Population II objects permeate the Milky Way Halo and Bulge, and many reside in the Disc. \\cite{1964ARA&amp;A...2...23P} and \\cite{1995CAS....27.....S} both provide excellent reviews of this class of variable stars. We refer the reader to Section 5 of \\cite{2006ARA&A..44...93S} for a review of optical ($V$-band) implementations of RR Lyrae stars as distance indicators.  It is well known that in optical wavebands the RR Lyrae period--luminosity relation is negligible (nearly no slope in the linear relation), but that at near- and mid-infrared wavebands the slope steepens and the relation is quite constrained. Figure 4 of \\cite{2013ApJ...776..135M} illustrates this phenomenon, and indicates an expected asymptote in the RRab period--luminosity relation slope of around $-2.4$ to $-2.8$. Additionally, since much of the Milky Way RR Lyrae population resides behind significant interstellar dust, the markedly reduced extinction at infrared wavelengths [demonstrated in Figure 8 of \\cite{2011ApJ...737...73F}] is often lower than photometric errors. Thus, at mid-infrared wavelengths, the need for dust extinction corrections to set mid-infrared luminosity vanishes.  In the present work we calibrate mid-infrared period--luminosity relations for RR Lyrae stars using the AllWISE Data Release of the {\\it Wide-field Infrared Survey Explorer} ({\\it WISE}\\,) and NEOWISE missions \\citep{2010AJ....140.1868W, 2011ApJ...731...53M}. Much of this work is a continuation of the methodology developed in \\cite{2011ApJ...738..185K}, which performed a similar analysis using the {\\it WISE} Preliminary Data Release. The earlier Preliminary Data Release included only the first 105 days of the {\\it WISE} mission, covering 57 per cent of the sky. The AllWISE Data Release (made public 2013 November 13) covers the entire sky and combines all the {\\it WISE} survey data from 2010 January 7 to 2011 February 1. By using this more comprehensive dataset the present analysis incorporates 129 RR Lyrae calibration stars, whereas  \\cite{2011ApJ...738..185K} made use of only 76 RR Lyrae period--luminosity relation calibrators. Furthermore, the longer temporal baseline and improved photometric accuracy of the AllWISE Data Release provides for more accurate mean-flux magnitude measurements of the calibration stars.  A similar investigation into the mid-infrared period--luminosity relations of RR Lyrae stars is presented in \\cite{2014arXiv1401.5523D}. The calibrations performed in \\cite{2014arXiv1401.5523D} use RR Lyrae stars in 15 Galactic globular clusters, with distances extending beyond 15~kpc.\\footnote{The calibration sample used in the present work, described in \\S\\ref{data_desc}, is comprised of nearby ($\\leqslant 2.5$ kpc) stars with well-detected {\\it WISE} light curves.} It is encouraging to note that the period--luminosity relation slopes reported in \\cite{2014arXiv1401.5523D}, modulo a slight metallically dependance, agree so well with the calibrations published in the present work.  The applications for RR Lyrae stars as distance indicators are diverse and fundamental to many areas of astronomical inquiry. They are found throughout the Milky Way Halo and Bulge, pervade the Disc, and are still shining in neighboring dwarf galaxies.  This paper is outlined as follows. We present a brief description of the {\\it WISE} and ancillary data in \\S\\ref{data_desc}. In \\S\\ref{methods} we review the employed light curve analysis methodology. In \\S\\ref{derivations} we present the derived period--luminosity relations, and in \\S\\ref{conclusions} we discuss the conclusions and future implications of this work.  \\section[]{Data Description}\\label{data_desc} {\\it WISE} provides imaging data in four mid-infrared wavebands: $W1$ centred at 3.4 $\\mu$m, $W2$ centred at 4.6 $\\mu$m, $W3$ centred at 12 $\\mu$m, and $W4$ centred at 22 $\\mu$m. Although the original {\\it WISE} mission was designed for static science goals, the orbit and survey strategy of the {\\it WISE} spacecraft (described in \\cite{2010AJ....140.1868W}) are conducive to recovering light curves of periodic variables with periods $\\lesssim 1.5$ day, which is well-matched to RR Lyrae variables.  In this analysis we make use of the most recent {\\it WISE} data release, AllWISE. The AllWISE Data Release (made public 2013 November 13) combines the 4-Band Cryogenic Survey (main {\\it WISE} mission covering the full sky 1.2 times from 2010 January 7 to 2010 August 6), the 3-Band Cryogenic survey (first three wavebands, 30 per cent of the sky from 2010 August 6 to 2010 September 29), and the NEOWISE post-cryogenic survey (first two wavebands, covering 70 per cent of the sky from 2010 September 20 to 2011 February 1). The individual photometry epochs were retrieved from the AllWISE Multiepoch Photometry Database.  As in \\cite{2011ApJ...738..185K} we employ the catalog of 144 relatively local ($\\leqslant 2.5$ kpc) RR Lyrae variables developed by \\cite{1998AA...330..515F}. Fifteen of these stars are excluded from our present analysis because they were not well-detected by {\\it WISE}. {\\it WISE} photometry data with any quality flags were rejected for the period--luminosity relation fits (most common cause was confusion with neighbouring sources, in part due to the $\\sim6.3$ arcsec PSF of {\\it WISE}\\,).  Also as in \\cite{2011ApJ...738..185K}, {\\it Hipparcos} photometry \\citep{1997ESASP1200.....P} is transformed into $V$-band \\citep{1998ApJ...508..844G}, corrected for dust extinction [using the line-of-sight extinction from \\cite{1998ApJ...500..525S} and the $R$ factor from \\cite{1975A&A....43..133S}], and combined with the \\cite{postHipp..book} $M_V$--[Fe/H] relation to yield prior distance moduli, $\\mu_{\\rm prior}$. Precise trigonometric parallax angles for four of the stars (RRLyr, UVOct, XZCyg, and SUDra; all of the RRab subclass) have been previously measured with the {\\it Hubble Space Telescope} and published in \\cite{2011AJ....142..187B}.\\footnote{RRc star RZCep also has an HST-measured parallax, but this star was rejected from the fit following the procedure described in \\S\\ref{methods}.} For these four stars the more precise, parallax-derived distance moduli are used in the period--luminosity relation fits. We note that the the distance moduli derived from the metallicity--luminosity relation for these four stars is in statistical agreement (within $2\\sigma$) with the parallax-derived distances.  A significant difference between the present work and that of \\cite{2011ApJ...738..185K} is that the two primary RR Lyrae subclasses (RRab and RRc stars) are now treated independently. Previously, in \\cite{2011ApJ...738..185K} all the RR Lyrae stars were fit together, which [as noted by \\cite{2013ApJ...776..135M}] is physically inappropriate because they follow different period--luminosity relations (RRab stars oscillate in the fundamental mode, whereas RRc stars do so in the first overtone). It is common practice to ``fundamentalize'' the periods of the RRc stars, as in \\cite{2004ApJ...610..269D}, and use them to supplement the RRab period--luminosity fit. In the present analysis we instead treat the two subclasses independently, so as to calibrate both the RRab and RRc period--luminosity relations. ",
        "conclusions": ""
    },
    "1402/1402.4480_arXiv.txt": {
        "abstract": "{The effect of the local environment on the evolution of dwarf spheroidal galaxies is poorly understood. We have undertaken a suite of simulations to investigate the tidal impact of the Milky Way on the chemodynamical evolution of dwarf spheroidals that resemble present day classical dwarfs using the SPH code GEAR. After simulating the models through a large parameter space of potential orbits the resulting properties are compared with observations from both a dynamical point of view, but also from the, often neglected, chemical point of view. In general, we find that tidal effects quench the star formation even inside gas-endowed dwarfs. {Such quenching, { may} produce the radial distribution of dwarf spheroidals from the orbits seen within large cosmological simulations.} We also find that the metallicity gradient within a dwarf is gradually erased through tidal interactions as stellar orbits move to higher radii. The model dwarfs also shift to higher $\\langle$[Fe/H]$\\rangle$/L ratios, but only when losing $\\gta$$20\\%$ of stellar mass. } ",
        "introduction": "Simulations of the prevailing cosmological {model} -- dark energy and cold dark matter, $\\Lambda$CDM -- predict that a copious amount of dark matter subhalos {will} survive to the present day. Within the Local Group these subhalos will be of a mass comparable to the halos that surround dwarf galaxies.  Although there is an underabundance of observed dwarf galaxies compared to the amount of subhalos predicted by $\\Lambda$CDM \\citep[e.g.][]{Klypin1999,Moore1999,Sawala2010}, numerous discoveries in the last decade \\citep{Willman2005,Willman2005a,Zucker2006,Belokurov2007,Irwin2007,Belokurov2008,Belokurov2009,Belokurov2010,Willman2011} and a greater understanding of baryonic processes has somewhat alleviated this \\citep{Bullock2010,Sawala2011}. All of these recently discovered dwarfs, and a majority of the previously known Local Group satellites are low luminosity, gas-deficient dwarf spheroidals (dSphs). Their low mass; proximity to the Galaxy (allowing individual stars, albeit only the brightest ones, to be resolved); and that they represent the final evolutionary state of many (though not all) dwarf galaxies \\citep{Grebel2003}; has made dSphs observational targets for studying a wide variety of processes related to galaxy formation.   Dynamical studies of dSphs have been undertaken by calculating the dynamical history of the Local Group dSphs and from numerical calculations determining the allowed phase-space of dSphs. Using the observed proper motions of dSphs \\citet{Lux2010} calculated the allowed perigalacticon and apogalacticon of seven Milky Way (MW) dSphs (and the Magellanic Clouds) by integrating the orbits backwards in time through a static potential. By calibrating the results to results from the Via Lactea II simulation they conclude that the MW's dSphs are likely to have formed before reionization. Importantly, the recorded dSphs with the lowest perigalacticon correspond to those with tidally induced inner bars \\citep{Lokas2012}. Similar tracing of orbits backwards in time has been undertaken by others for individual galaxies \\citep[e.g.][]{Pasetto2011} finding similar results. In all cases the uncertainty in the proper motion and the mass accretion history of the MW limits the accuracy of the resolved orbit.  In contrast to determining the orbital history of the dwarfs by integrating the observed motion backwards in time, \\citet{Rocha2012} compared the present day {kinematical} properties of dSphs with the Via Lactea II simulation to determine a likely accretion time for these dwarfs. This accretion time (aka the infall time, the time at which a satellite galaxy first crosses the virial radius of its host) is compared to the star formation history of the dwarf galaxies. Here, the ultrafaint dSphs are found to have likely ceased star formation before infall. For some classic dSphs there is an overlap between possible accretion times and dramatic increases in the star formation before being quenched within approximately a Gyr of infall. That such overlaps exist for a few of the classical dwarfs (although not all, Carina for instance shows a burst of star formation post any likely infall time) suggests the environment strongly influences the star formation rate of dSphs, potentially being responsible for the quenching of dwarf galaxies \\citep{Rocha2012}.  Complementing the orbital history calculations of dSphs, some authors have examined the effect of interactions between the MW and the orbiting dSph. Dynamically the two major components of these interactions are tidal forces arising from the dSphs interaction with the gravitational field of the MW, and gas stripping due to the interaction of the dSph with the predicted hot gas halo of the MW \\citep{Blitz2000,Grcevich2009,Miller2013,Gatto2013}.  Examining the tidal forces acting upon simulated dwarfs, \\citet{Penarrubia2008} found that many dwarf galaxies will be able to retain a stellar density profile of similar shape to their original profile, even when retaining only a small fraction ($\\sim$$1\\%$) of their stars. Due to the \"cuspy\" nature of dark matter protecting it from stripping, such dwarfs have evolved elevated dynamical to luminous matter ratios, similar to that of dSphs. Consequently, the maximum circular velocity of the dwarf was found to drop in line with the decreasing stellar mass. This dependence is not seen inside the local dSphs and hence \\citet{Penarrubia2008} found it unlikely that dSphs are tidally evolved remnants from more massive dwarf galaxies.  {\\citet{Kazantzidis2013} explored the effect of the dark matter potential on this stripping and found that dwarfs with a less \"cuspy\" profile are more likely to be transformed into systems reminiscent of dSphs. In particular for these core-like dwarfs, such transformation can take place within one pericentric passage and dwarfs on near circular orbits and at large perigalacticons are able to be transformed. When considering which dwarfs are destroyed, they find that dSphs are more likely to be found on circular orbits or at higher perigalacticons. With eccentric, low perigalacticon dwarfs ending up as tidal streams.}  \\citet{Mayer2006} examined both tidal and ram pressure stripping through high resolution {\\em N}-body+SPH simulations of gas rich dwarfs (sans stellar feedback or chemical evolution) within a MW like halo. Tidal stirring of the dwarfs produces objects with kinematics resembling dSph while ram pressure stripping removes the gas at low perigalacticons. Such processes were sufficient to strip low mass dwarfs with low perigalacticons and with ultraviolet heating from the extragalactic UV field (or what may be expected from stellar feedback), stripping of more massive dwarfs at similar perigalacticons. That ram pressure stripping is sufficient to remove this gas at low perigalacticons agrees with analytic predictions for dwarf galaxies \\citep{Mori2000}. It is important to note however, that SPH simulations, such as those undertaken by \\citet{Mayer2006} suffered from numerical problems \\citep[see][]{Agertz2007} due to the artificial viscosity present in standard SPH formulations which enhanced the survivability of gaseous clouds against stripping and predictions concerning ram pressure stripping may suffer from some inaccuracy as a consequence of this.  Although tidal and ram pressure stripping are sufficiently powerful to strip galaxies close in (perigalacticons$\\lta50$~kpc) to the host galaxies, the gas deficiency of dwarfs continues well past $200$~kpc \\citep{Grcevich2009}. At these large distances, stellar feedback may provide sufficient heating to allow gas to be readily stripped even by the sparse {hot gas} halo \\citep{Nichols2011,Nichols2013}. Understanding how the star formation history of dwarf galaxies is therefore immensely important in understanding how they evolve over cosmic time.  Parallel to studies of dynamical evolution, much work has been conducted trying to understand the chemical evolution of dwarf galaxies. With evolutionary histories that often contains periods of heightened star formation after the initial burst \\citep[e.g.][]{Tolstoy2009,Lee2009} and large metal loss which has been proposed analytically to be expelled through supernova driven winds \\citep{Kirby2011b}. These heightened periods, seen in both isolated and tidally interacting dwarfs \\citep[{although such dwarfs are potentially more massive than the dSphs today as seen by their ability to retain gas}]{Brosch2004} may provide an environment where metals may quickly pollute the ISM before the next period of star formation commences, however, how quickly the metals are able to be redistributed and the quantity {lost during the burst} remain open questions.    We have previously examined how the star formation history imprints on the chemical abundances of dwarf galaxies through chemodynamical processes in isolated dwarfs using the TreeSPH code GEAR \\citep{Revaz2009,Revaz2012}. The code fully solves the gravitational and hydrodynamics in addition, it accounts for star formation processes, chemical enrichments, cooling and supernova feedback. \\citet{Revaz2009,Revaz2012} obtained generic models that described the great variety of dSphs observed in the local universe. The mass ranges and star formation histories of these generic models are compatible the with observed dSphs like Carina, Sextans, Sculptor and Fornax. {However, star formation continued on past the initial burst, preventing the development of dwarfs with Sextans or Sculptor like star formation histories which display a single, terminated, period of star formation \\citep{LeeMG2009,deBoer2012}. In order to fully match the present day chemical evolution, this star formation had to be artificially terminated for Sextans and Sculptor. As these dwarfs relied only on internal processes, the requirement of artificial termination is in line with previous work, with \\citet{Sawala2010} demonstrating that internal feedback is not sufficient to reproduce the full range of observed galaxies. In particular to reproduce dwarfs that only have a short burst of star formation initially before it infall onto a larger galaxy, additional UV heating (such as from the extragalactic UV field) {may be} necessary \\citep[see also e.g.][]{Ricotti2005,Valcke2008}.}  {Even with the artificial termination in star formation history reproducing the observed metallicity, mass and luminosity}, the model galaxies retained large quantities of gas ($\\gta10^{7}~M_{\\sun}$) to the present day, in conflict with the gas deficiency observed {in satellite galaxies within the Local Group} \\citep{Grcevich2009}. However, the closest dwarf galaxies (and subsequently the ones we are best able to study) are -- by virtue of their proximity -- not isolated. How the interaction between a low mass galaxy and an MW sized galaxy affects the chemical abundances inside the dwarf galaxy is currently unknown, with no previous studies exploring this. Importantly, is it possible to have the gas removed from dwarfs just through this interaction without greatly modifying the chemical or star formation properties of a dwarf, or are calculations of dwarfs in isolation unsuitable for the nearest dwarfs{?} Here we examine whether tidal interactions between model dSph and the MW greatly alter the chemodynamical properties of the dSph and what this may reveal about the progenitor properties of some of the classical dwarf spheroidals. In \\S\\ref{sec:IC} we describe {the GEAR code}, model setup and initial conditions of the model dwarfs, the impact of tidal forces are discussed in \\S\\ref{sec:tide}, with concluding remarks in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}   We have undertaken a series of simulations to examine how the chemodynamical properties of dwarf spheroidals are altered through tidal interactions. That major chemodynamical changes, excluding those arising from star formation quench, were only minor, although potentially noticeable, in undisrupted dSphs suggests that the isolation models provide a good base to understanding how dSphs have evolved through cosmic time.  As our base models, which was evolved in isolation were chosen to closely match the observed properties of dSphs when evolved in isolation \\citep{Revaz2012}, na\\\"\\i{}ve comparisons with observations will bias towards those models with the smallest tidal impacts. However, as nearly all dSphs that experienced drastic changes in observable properties such as metallicity or luminosity (see also Appendix B) became noticeably disrupted, such a bias is likely to be small when considering at least the classical dSphs.  {The quenching of star formation seems to occur in more massive dwarfs, and biased towards those undergoing star formation near the perigalacticon. If the stellar feedback is weaker than proposed, such an effect may be reduced to the outer edges of a dwarf, or disappear entirely. As this effect was somewhat mass dependent, the dwarfs most likely to experience such a quenching near perigalacticon are those which underwent star formation before infall and hence are able to blow gas away from the central star forming region. }  { Due to the large parameter space that satellite galaxies inhabit and computational restraints there are limitations to the modelling able to be undertaken. The choice of a cored profile arises from the optimisation of the best fit models in isolation, with the initial conditions for dSphs in an NFW profile so far unknown. Replacing the density profile, without changing any other properties, does allow a dwarf slightly more resilience to tidal stripping. But such a dwarf also over produces stars for the same baryon/dark-matter ratio, and consequently a different chemical history, suggesting that a lower mass NFW profile may be appropriate. In this case, the tidal effects still lead to disruption and for the masses considered, a quenching of star formation, however, the central region may maintain more consistency than the cored profile examined here.}  Our simulations were also undertaken with a mass of the Milky Way halo on the low tail of the distribution proposed in the literature \\citep{Boylan-Kolchin2011,Kafle2012} although still in good agreement with observations. Using a more massive halo is likely to extend its sphere of influence to disrupt dSphs even further out than that proposed here. {Such a result would be in contrast to that if run inside a growing cosmological simulation. Due to the lower mass at early times, the tidal forces could be expected to be lower, however, as the evolution of orbits within a growing halo are heavily phase dependent it is difficult to say definitively.}  Our investigation also ignored the ram pressure effects of a hot halo, such a hot halo is likely to act in synergy with tidal stripping to remove the low density gas from the dSphs potential, allowing the dwarf to then expel any gas that survives (as shown in Fig. \\ref{fig:gasexpulsion} but occurring over a much larger orbital range). We also explored only a small portion of the parameter space of dwarf progenitors, more massive dSphs may be expected to better survive tidal interactions while least massive will be more easily disrupted.    We find that\\begin{itemize} \\item tidal forces during one perigalacticon passage are sufficient to quench star formation even in gas endowed dSphs, \\item the metallicity gradient is affected by the number of passages, with more tidal interactions resulting in a lower metallicity gradient, \\item dSphs are generally inefficient at the recycling of metal from supernovae and will in all cases only have a few percent of metals recycled no matter the state of remaining gas. \\end{itemize} The quenching of star formation is due to synergy between supernovae and tidal forces keeping gas at sufficiently low densities to prevent star formation and we do not see any gas compression in these dSphs which may allow further star formation to occur during these interactions. Such quenching suppresses the latest generation of stars, an effect which imprints into the $\\alpha$/Fe abundance ratios within the dSphs. The halting of star formation supports the suggestion that the radial distribution of dSphs arises from tidal interactions \\citep{Slater2013}.  A portion, approximately half, of the lower metallicity gradient was also due to the quenched star formation, however, tidal interactions do partially erase metallicity gradients through purely tidal effects.  Such tidal interactions also shift dSphs to a higher $\\langle$[Fe/H]$\\rangle$/L ratio, allowing a greater range of the observed distribution covered than dSphs that have only evolved in isolation, albeit only when$\\gta$$20\\%$ of stars are lost. The dSphs which undergo two passages are able to lose $30\\%$ of their initial gas mass purely to tidal forces and internal processes, while which only undergo one passage can only lose $5\\%$ the same way, with dSphs that lose more tending to undergo too much disruption to be recognisable as classical dSphs today. Such small gas losses, when nearly all dSphs are \\ion{H}{i} deficient, indicates that other processes are necessary to strip the majority of a dSphs initial gas. Ram pressure is the most likely process for this stripping, with it able to remove the gas from a dwarf rapidly \\citep[e.g.][]{Gatto2013}. Dwarf winds may also contribute, although such winds are not needed to instantly remove metals from the dSphs and would have to work in conjunction with ram pressure to produce the observed radial distribution of gas-rich/gas-deficient dwarfs \\citep[see][]{Grcevich2009}."
    },
    "1402/1402.6344_arXiv.txt": {
        "abstract": "Deep late-time X-ray observations of the relativistic, engine-driven, type Ic SN\\,2012ap allow us to probe the nearby environment of the explosion and reveal the unique properties of relativistic SNe. We find that on a local scale of $\\sim0.01\\,\\rm{pc}$ the environment was shaped directly by the evolution of the progenitor star with a pre-explosion mass-loss rate $\\dot M <5 \\times 10^{-6}\\,\\rm{M_{\\sun}\\,yr^{-1}}$, in line with GRBs and the other relativistic SN2009bb. Like sub-energetic GRBs, SN\\,2012ap is characterized by a bright radio emission and evidence for mildly relativistic ejecta. However, its late time ($\\delta t\\approx 20\\,\\rm{d}$) X-ray emission is $\\sim100$ times fainter than the faintest sub-energetic GRB at the same epoch, with no evidence for late-time central engine activity. These results support  theoretical proposals that link relativistic SNe like 2009bb and 2012ap with the weakest observed engine-driven explosions, where the jet barely fails to breakout. Furthermore, our observations demonstrate that the difference between relativistic SNe and sub-energetic GRBs is intrinsic and not due to line-of-sight effects. This phenomenology can either be due to an intrinsically shorter-lived engine or to a more extended progenitor in relativistic SNe. ",
        "introduction": "\\label{Sec:Intro}  The vast majority of supernova explosions (SNe) arising from hydrogen and helium stripped progenitors (i.e. type Ic SNe, see \\citealt{Filippenko97} for the spectral classification of SNe) can be explained by the hydrodynamical collapse of the massive progenitor star (e.g. \\citealt{Tan01}). In a very limited percentage of cases ($\\lesssim 1$\\%, \\citealt{Berger03a,Coward05}, \\citealt{Guetta07}, \\citealt{Soderberg10}), the explosion is instead powered by an engine able to accelerate a tiny portion of the ejecta with typical mass\\footnote{For the relativistic SN\\,2009bb  $M\\approx 10^{-2}\\,\\rm{M_{\\sun}}$ \\citep{Chakraborti11}.} $M\\approx 10^{-5}-10^{-6}\\,M_{\\sun}$ to velocities $v\\gtrsim0.6\\,c$. Engine-driven explosions (E-SNe hereafter) are thus uncommon. Furthermore, only a small fraction of E-SNe ( $\\lesssim10\\%$, e.g. \\citealt{Soderberg06}, see also \\citealt{Cobb06,Pian06, Liang07,Virgili09}) harbor a fully relativistic jet and give origin to Gamma-Ray Bursts (GRBs). The peculiar circumstances that cause a hydrogen-stripped, massive progenitor star to produce a relativistic jet at the time of the collapse are still not fully understood. High angular momentum seems to be a key ingredient (e.g. \\citealt{MacFadyen99}, \\citealt{MacFadyen01}, \\citealt{Woosley06}, \\citealt{Dessart08}).  E-SNe have historically been detected through their prompt X-ray and $\\gamma$-ray emission produced by energy dissipation within the jet (ordinary GRBs) and by the SN shock break out (which is relevant at least for some sub-energetic -sub-E- GRBs, e.g. \\citealt{Kulkarni98,Matzner99,MacFadyen01,Tan01,Campana06, Wang07,Waxman07,Katz10,Bromberg11,Nakar12}). More recently, two E-SNe have been discovered through their bright later-time radio emission (i.e. the relativistic SNe 2009bb and 2012ap, \\citealt{Soderberg10,Bietenholz10,Chakraborti11b}; \\citealt{Chakraborti14}, hereafter C14). Here we specifically ask the questions: what is the nature of relativistic SNe and what is their connection with the other classes of E-SNe known so far?   The two known relativistic SNe 2009bb and 2012ap share with sub-E GRBs  evidence for mildly relativistic ejecta powering a bright radio emission (\\citealt{Soderberg06,Soderberg10,Bietenholz10}, C14) and a very energetic optical explosion with $E_{\\rm{k}}\\sim10^{52}\\,\\rm{erg}$ coupled to material moving at $v\\sim$ a few $10^4\\,\\rm{km\\,s^{-1}}$ (\\citealt{Pignata11,Milisavljevic14b}). These properties dynamically distinguish relativistic SNe and sub-E GRBs from ordinary SNe, and put these explosions between the highly-relativistic, collimated GRBs and the more common type Ic SNe.  On the theoretical side, state-of-the art simulations of jet-driven stellar explosions (e.g. \\citealt{Lazzati12}) associate classic GRBs with fully-developed, highly-relativistic jets and suggest that sub-E GRBs likely represent the cases where the jet is just barely able to pierce through the stellar envelope (\\citealt{Bromberg11}, \\citealt{Nakar12}). In particular, it was suggested by \\citealt{Lazzati12} that a different life-time of the central engine might be able to explain the entire zoo of E-SNe (i.e. relativistic SNe, sub-E GRBs and ordinary GRBs). However, it is unclear if relativistic SNe represent a new class of explosions, or if, instead, they are the equivalent of sub-E GRBs for which we missed the high-energy trigger because of line-of-sight effects or incomplete coverage of the $\\gamma$-ray satellites (see e.g. the discussion for SN\\,2009bb in \\citealt{Soderberg10}). This still-open-question motivates this study.  We present late-time deep X-ray observations of the relativistic SN\\,2012ap. These observations allow us to identify for the first time a distinctive property of relativistic SNe that clearly sets them apart from all the other known engine-driven explosions. We find that relativistic SNe are characterized by a significantly fainter X-ray emission at late times ($t\\sim 20$ d), even compared to sub-E GRBs (SN\\,2012ap is $\\sim100$ times fainter than the faintest sub-E GRB at the same epoch), and show no evidence for an excess of X-ray radiation that has been recently reported for the sub-E GRBs 060218 (\\citealt{Soderberg06, Fan06}) and 100316D (\\citealt{Margutti13b}).  We describe our observations in Sec. \\ref{Sec:Obs} and constrain the progenitor mass-loss rate in Sec. \\ref{Sec:massloss}. Finally we put SN\\,2012ap in the context of engine-driven explosions in Sec. \\ref{Sec:CE} and \\ref{Sec:discussion} and discuss how our findings clearly suggest that relativistic SNe constitute a separate class of engine-driven explosions with intrinsic differences with respect to sub-E GRBs. Conclusions are drawn in Sec. \\ref{Sec:conclusions}.  Uncertainties are quoted at $1\\sigma$ confidence level, unless otherwise noted. We employ standard cosmology with $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73$, and $\\Omega_{\\rm M}=0.27$. Throughout the paper we use 2012 February 5th as the explosion date of SN\\,2012ap, as inferred by \\cite{Milisavljevic14b} (M14, hereafter) from extensive optical observations. Following  C14 we assume a distance of 40 Mpc \\citep{Springob07,Springob09}. A detailed discussion of the optical and radio properties of SN\\,2012ap can be found in M14 and C14, respectively. Finally we note that sub-E GRBs are also called low-luminosity GRBs in the literature (see e.g. \\citealt{Bromberg11}). However the physical parameter that is relevant to your analysis is the (modest) kinetic energy of their fastest ejecta, which is surely related to the low-luminosity of their prompt $\\gamma$-ray emission. For this reason we will refer to this class as sub-E GRBs. ",
        "conclusions": "\\label{Sec:conclusions} The class of engine-driven SNe (E-SNe) collects a rare variety of SN explosions ($\\lesssim1\\%$ of type Ic SNe) and includes relativistic SNe, sub-E GRBs and ordinary GRBs. E-SNe are characterized by a significantly shallower kinetic energy profile of the explosion ejecta than  expected in the case of a pure hydrodynamic collapse of the progenitor star (Fig. \\ref{Fig:EkGammaBeta}), indicating that E-SNe are able to accelerate a tiny but important fraction of their ejecta to higher velocities ($v\\gtrsim 0.6\\,c$). E-SNe otherwise show a diverse phenomenology.  Relativistic SNe and sub-E GRBs share a bright radio emission and evidence for mildly relativistic ejecta that clearly set them apart from ordinary SNe Ic (non-relativistic). The thermal properties of relativistic SNe are also analogous to the very energetic, fast expanding SNe associated with sub-E GRBs and ordinary GRBs (\\citealt{Soderberg10}, M14). \\emph{However, a distinctive property of relativistic SNe is their significantly fainter X-ray emission (Fig. \\ref{Fig:Xray}) that implies the lack of a luminous X-ray component arising from the central engine activity at late times ($\\delta t\\sim 20\\,\\rm{d}$, Fig. \\ref{Fig:Radio2Xray}).} With $E_{\\rm{\\gamma,iso}}<10^{47}\\,\\rm{erg}$, the prompt $\\gamma$-ray energy released by the relativistic SN\\,2012ap is also considerably below the level of sub-E GRBs (Fig. \\ref{Fig:LxEiso}).\\footnote{For the other relativistic SN\\,2009bb the observational limit on the promptly released  $E_{\\rm{\\gamma,iso}}$ is unfortunately not as constraining (Fig. \\ref{Fig:LxEiso}).} The higher metallicity of the environment and the conspicuous presence of helium in the ejecta of the two known relativistic SNe also sets them apart from sub-E GRBs and GRBs.  These findings call for some crucial diversity in the properties of the engines and/or of the progenitors of relativistic SNe and sub-E GRBs (and ordinary GRBs as well). We showed that the observations of relativistic SNe are consistent with the picture of jet-driven explosions where the jet just barely fails to breakout from the progenitor star. This scenario naturally explains (i) the lack of evidence for central engine activity at late times and (ii) the deep limit to the promptly released $\\gamma$-ray energy.  The failed jet breakout might be due to an intrinsically short-lived engine (but same progenitor properties) or to a different progenitor structure between relativistic SNe and sub-E GRBs. At the time of writing, with only two relativistic SNe discovered so far (through targeted optical surveys), observations do not allow us to distinguish between these two scenarios. A significantly larger sample of relativistic SNe found through \\emph{untargeted} SN searches in the optical and radio band is clearly needed to deeply understand their connection to GRBs, build a complete picture of E-SNe and constrain which unique property differentiates failed breakouts from successful, fully relativistic jets. The Large Synoptic Survey Telescope (LSST, \\citealt{Ivezic08}) is expected to discover hundreds of SNe Ic every year, thus in principle providing the significantly larger sample of E-SNe that is needed to deeply understand the connection between relativistic SNe and GRBs.  However,  as we demonstrate here,  coordinated radio \\emph{and X-ray} follow up is essential to identify E-SNe from their ordinary counterparts and determine the properties of the engines that power their explosion."
    },
    "1402/1402.4800_arXiv.txt": {
        "abstract": "The observational signatures of multi-field inflation will generally evolve as the Universe reheats. We introduce a general analytic formalism for tracking this evolution through perturbative reheating, applicable to two field models with arbitrary separable potentials. The various transitions, including the onset of scalar field oscillations and the reheating of each field, can happen in different orders and on arbitrary hypersurfaces. The effective equations of state of the oscillating fields are also arbitrary. Nevertheless, our results are surprisingly simple. Our formalism encapsulates and generalises a huge range of previous calculations including two-field inflation, spectator models, the inhomogeneous end of inflation scenario and numerous generalised curvaton scenarios. ",
        "introduction": "The presence of isocurvature modes in the early Universe has a profound influence on the evolution of cosmic observables. During the slow-roll inflationary phase, this effect has been tracked analytically by Vernizzi and Wands~\\cite{Vernizzi:2006ve}, extending earlier work~\\cite{GarciaBellido:1995qq}, for two-field models with arbitrary sum-separable potentials. Their results apply to models where reheating occurs suddenly, such as two-field Hybrid inflation where inflation is terminated by the destabilisation of a third heavy field~\\cite{Alabidi:2006wa, Alabidi:2006hg, Sasaki:2008uc, Byrnes:2008zy}. In such cases, the isocurvature modes are assumed to decay almost instantaneously due to the onset of a phase of thermal equilibrium. It is possible that this process of sudden reheating may occur on a non-uniform density hypersurface. This is known as the Inhomogeneous End of Inflation scenario~\\cite{Bernardeau:2002jf,Lyth:2005qk,Salem:2005nd,Alabidi:2006wa,Bernardeau:2007xi,Sasaki:2008uc,Naruko:2008sq,Huang:2009vk} and has been recently summarised and constrained by Planck data in ref.~\\cite{Elliston:2013afa}. Since the method of Vernizzi and Wands~\\cite{Vernizzi:2006ve} tracks perturbations until the end of inflation, their work is equally applicable to any two-field separable potential where observable statistics attain constant conserved values before the end of inflation. This happens when the isocurvature modes decay during inflation and the system reaches an `adiabatic limit'~\\cite{Meyers:2010rg,Elliston:2011dr}.  More generally, isocurvature modes may be present after inflation, and inflation may not end suddenly. This will mean that both the inflationary and the post-inflationary phases will cause the cosmic observables to evolve. As emphasised in Elliston et al.~\\cite{Elliston:2011dr}, robust predictions can only then be generated if the primordial perturbations are tracked through the post-inflationary phases, either until isocurvature modes do finally decay, or until the time at which the system is observed. In particular, it is of vital importance to consider how cosmic observables may be sensitive to the process by which the Universe reheats. In addition to the sudden non-perturbative reheating mechanism present in Hybrid inflation, perturbative reheating can also be considered. In this case an oscillating field dissipates its energy more gradually, often significantly after the end of inflation. Therefore any analytic calculation for a model exhibiting this behaviour must give due consideration to the post-inflationary phases. The most prominent example of the effect of perturbative reheating on the generation of primordial perturbations is the `curvaton' scenario~\\cite{Lyth:2001nq,Moroi:2001ct,Enqvist:2001zp,Sasaki:2006kq}.  A third method to reheat the Universe is preheating, driven by explosive resonant particle production. Like the end of Hybrid inflation, this form of reheating is sudden and non-perturbative. If the field into which energy is being transferred is light, it can have a significant effect on the evolution of perturbations~\\cite{Bassett:1999cg,Finelli:2000ya,Chambers:2007se, Bond:2009xx}. Since such effects must be studied using lattice simulations, however, we shall not consider this behaviour in our analytic work, restricting our attention instead to perturbative reheating.  The goal of this paper is to augment the work of Vernizzi and Wands~\\cite{Vernizzi:2006ve} to derive a simple and usable set of analytic formulae that allow us to understand how perturbative reheating modifies cosmic observables. Our method can be applied to any two-field sum-separable model. We draw upon and generalise many different calculations that have been tailored to more specific scenarios; these scenarios may then be recovered as particular limits of our work. In particular, we account for the Modulated Reheating scenario~\\cite{Dvali:2003em,Zaldarriaga:2003my,Kofman:2003nx,Vernizzi:2003vs,Bernardeau:2004zz,Ichikawa:2008ne,Suyama:2007bg} where the hypersurface on which a particular field reheats is directly dependent on the value of an additional light field, and for a multitude of possible generalisations to the standard curvaton scenario. These include modulation of the onset of curvaton oscillations~\\cite{Kawasaki:2011pd, Kawasaki:2012gg}, modulation of the reheating hypersurface for the curvaton~\\cite{Enomoto:2012uy,Assadullahi:2013ey,Langlois:2013dh,Assadullahi:2012yi,Enomoto:2013qf}, and modulation of the reheating hypersurface for the inflaton~\\cite{Choi:2012te}. See refs.~\\cite{Elliston:2013afa,Vernizzi:2003vs,Alabidi:2010ba} for other works drawing links between these scenarios.  Our work follows a considerable amount of recent effort in this field. Numerical studies~\\cite{Elliston:2011dr,Leung:2012ve} have demonstrated that altering the decay widths of one or both fields causes the local shape bispectrum parameter $\\fnl$ to evolve. Subsequently, Meyers and Tarrant~\\cite{Meyers:2013gua} demonstrated that this evolution was caused by the same physical mechanism operating behind the curvaton scenario; the curvature perturbation is modified by the relative redshifting of two components with different equations of state. They showed this by providing general analytic formulae that describe the evolution of perturbations during the perturbative reheating phase. The effect of perturbative reheating has also been discussed in the context of particular models in refs.~\\cite{Enomoto:2013bga,Huston:2013kgl}.  The formalism that we develop in this paper is more general than that of Meyers and Tarrant~\\cite{Meyers:2013gua}. We allow for reheating to occur on arbitrary hypersurfaces, allowing us to consider the modulated reheating scenarios discussed above. Since there is no agreed analytic method for fixing the hypersurfaces of oscillation onset, we also leave these hypersurfaces arbitrary. This ensures that our formalism, instead of being limited by the lack of a well-defined oscillation criterion, may be used as a testing-ground for choosing between methods of defining the oscillation hypersurfaces. The explicit nature of our formalism also makes it very simple to interpret and to obtain standard limiting cases.  We split the evolution into four phases, labelled A, B, C and D, which are summarised in figure \\ref{fig:summary}. At any particular time, the two components involved will be described either as scalar fields undergoing generalised slow-roll, or as fluids with a constant equation of state. Phase A contains two slow-roll fields and is therefore a reiteration of the results of ref.~\\cite{Vernizzi:2006ve}. The four-phase picture we employ is a choice which covers the vast majority of models encountered in the literature, but it is not exhaustive. However, the basic tools that we develop can be assembled by the user to tackle any similar problem even if it falls outside of our picture. \\\\  \\noindent {\\bf Paper structure}: In \\S\\ref{sec:background_theory} we introduce the necessary background theory. \\S\\ref{sec:deltaN} then shows how the curvature perturbation $\\zeta$ may be computed in terms of horizon exit field perturbations and provides the principle result of this paper. \\S\\ref{sec:ns} then provides a general discussion regarding predictions for the scalar spectral index and the tensor to scalar ratio. \\S\\ref{sec:second_order} provides the ingredients for deriving predictions at second order, which may be assembled to suit specific problems as we show in \\S\\ref{sec:models}. We conclude in \\S\\ref{sec:conclusions} and some supporting formulae are given in the appendix. ",
        "conclusions": "\\label{sec:conclusions}  This paper provides an analytic formalism for showing how perturbative reheating modifies the predictions of two-field slow-roll inflation. This is therefore an extension of the inflationary work carried out by Vernizzi and Wands~\\cite{Vernizzi:2006ve}. Our calculations are made possible by replacing the oscillating fields with effective fluids.  The hypersurfaces where the two fields begin to oscillate are arbitrary, as are the hypersurfaces where the two fields reheat. This allows us to include all of the effects of modulation throughout. In addition, we do not require that the fields reheat in any particular order. Despite this generality, the linear order contribution to $\\zeta$ in eq.~\\eqref{eq:dN_3} is reasonably simple. This linear order result is the principle result of this paper, since all higher order results follow by differentiation. We provide the ingredients needed to compute the second order results, which may be assembled to suit a given problem. Our explicit results are specific to the four phases illustrated in figure~\\ref{fig:summary} which reduces to a large range of existing scenarios in different limits. The calculations that define our formalism are more general, however, and can easily be adapted to models that do not fit this four-phase picture, such as the inflating curvaton.  There is a clear physical motivation to consider modulation of the reheating hypersurfaces, since this allows us to fully account for the physical effects of modulated reheating or an inhomogeneous end of inflation. The main motivation to allow the fields to begin oscillations on an arbitrary hypersurface is simply that we do not currently have a robust prescription for what this hypersurface should be. Earlier work by Kawasaki et al.~\\cite{Kawasaki:2012gg} demonstrated that this can be different from a surface of uniform density and they provided a simple prescription that works well for spectator models---where one field dominates the energy density. We then tested this prescription outside of the spectator regime and found that it produces spurious effects on the predictions for observables. Further work is therefore required to discern the analytic conditions that define the onset of oscillations. Since our formalism is agnostic about such analytic conditions, it provides a suitable test bed for studying different suggestions.  We have used our formalism to discuss predictions for a model where reheating is modulated by a fluid rather than a field. The conclusions show that a large value of $\\fnl$ is easily produced if the hypersurface has the correct geometry. This is consistent with the results of Elliston et al.~\\cite{Elliston:2013afa} where it was shown that the geometry of the reheating hypersurface is significantly constrained in the light of Planck bispectrum data. We therefore expect that this scenario may be similarly constrained.  Our work has important consequences for earlier numerical work. In particular, Leung et al.~\\cite{Leung:2012ve} discussed the evolution of $\\fnl$ as one varies the $\\Gamma$ parameters of perturbative reheating. Like Meyers and Tarrant~\\cite{Meyers:2013gua}, our work demonstrates that these numerical findings can be neatly explained by considering the relative redshifting of the two fluids involved."
    },
    "1402/1402.0338_arXiv.txt": {
        "abstract": "{The expansion of coronal loops in the transition region may considerably influence the diagnostics of the plasma emission measure. The cross sectional area of the loops is expected to depend on the temperature and pressure, and might be sensitive to the heating rate.} {The approach here is to study the area response to slow changes in the coronal heating rate, and check the current interpretation in terms of steady heating models.} {We study the area response with a time-dependent 2D MHD loop model, including the description of the expanding magnetic field, coronal heating and losses by thermal conduction and radiation from optically thin plasma. We run a simulation for a loop 50 Mm long and quasi-statically heated to about 4 MK.} {We find that the area can change substantially with the quasi-steady heating rate, e.g. by $\\sim 40$\\% at 0.5 MK as the loop temperature varies between 1 and 4 MK, and, therefore, affects the interpretation of DEM(T) curves.} {} ",
        "introduction": "\\label{intro}  An important challenge in trying to understand the structure and heating of solar active regions is to explain the relative amounts of material at different temperatures. This is typically expressed in terms of the differential emission measure distribution, DEM($T$). Models that use steady heating and a flux tube geometry that slowly expands with height are unable to reproduce the observed distributions. However, realistic flux tubes are expected to expand very rapidly with height between the chromosphere and corona, where the plasma beta ($\\beta = 8 \\pi P / B^2$) changes from being much greater than to much less than unity. The thin transition region where temperatures increase from chromospheric to coronal values occurs somewhere within this rapidly expanding throat. Thus, we expect the cross sectional area of flux tubes to be a strong function of temperature. If the dependence has the right form, steady heating models may in fact be able to explain the observed DEM(T). This was demonstrated by \\cite{Warren_al_2010_b} for the core of an active region (the generally hot central part of the active region that connects opposite polarity moss regions, which are the transition region footpoints of the hot loops).  A potential difficulty with this picture has been pointed out by \\cite{Tripathi2010a}. The height of the transition region depends on the coronal pressure and so moves up and down in response to changes in the coronal heating rate, either fast or slow. If the heating rate increases, so too does the pressure, and the transition region is forced downward. If the heating rate decreases, some of the overlying pressure is relieved, which allows the transition region to move upward. For gradual heating variations, the amount of vertical displacement is given approximately by  \\begin{equation} \\Delta z \\approx -\\frac{7}{6} H_g \\ln \\left(\\frac{Q}{Q_0}\\right) , \\label{eq:deltaz} \\end{equation}  \\noindent where $H_g$ is the gravitational scale height of the chromosphere (approximately 500 km), $Q$ is the new heating rate, and $Q_0$ is the original heating rate. We have made use of the fact that $Q \\propto p^{7/6}$ for quasi-static evolution, where $p$ is the pressure. See \\cite{Klimchuk_2006} for a similar expression appropriate to impulsive heating. A heating rate increase of a factor of 100, which increases the loop apex temperature by a factor of 3.7, causes the transition region to be displaced downward by approximately 2700 km. This is not small compared to the length scale of expansion in the flux tube throat \\citep{Gabriel_1976,Athay1981a}. Thus, the cross sectional area at any given temperature $A(T)$ could change dramatically. This would alter DEM($T$), so that a model is no longer compatible with observations. The very sensitive dependence of the $A(T)$ relationship on the heating rate suggests that too much fine tuning is required for a steady heating model to have widespread applicability.  The last statement involves a big caveat, however. It assumes that flux tubes remain rigid as the heating and pressure change. This is not the case in a real MHD system. With increasing pressure, a flux tube will expand laterally at the same time that the transition region is pushed downward. Hence, the area at a given temperature might remain constant, whereas it would decrease in a rigid tube. It is possible that the $A(T)$ relationship in the transition region could remain largely unchanged. (We note, however, that any lateral expansion would be even more pronounced in the corona, so the emission measure relationship between the transition region and corona would then change.) The purpose of our study is to determine how $A(T)$ varies in response to changes in the coronal heating rate, especially slow changes corresponding to quasi-static evolution. If there is a strong dependence on the heating rate, then the steady heating interpretation of DEM($T$) should be reexamined.   In this paper we present a 2D-MHD loop model that naturally accounts for loop expansion  through the transition region. We have used this model to study the plasma response to the heating into and around the moss regions. \\rev{At variance with 1D models having a specified expanding cross section \\citep{Emslie1992a,Mikic2013a}, we use a full MHD description, i.e., the full set of MHD equations is solved altogether in a 2D spatial domain, that allows us to describe a beta-changing system and to have feedbacks between plasma and magnetic field in the critical region, i.e. around the transition region. As a consequence, and at variance with previous works, in our model the loop cross-section area is also a function of time and changes as the heating changes. Our 2D cylindrical coordinates make the description realistic for a 3D structure with cylindrical symmetry, that is a good approximation for a coronal loop (but see \\citealt{Malanushenko2013a} for deviations from circular loop cross-sections). We have focussed on the expansion around the transition region, not in the whole loop \\citep{Peter2012a}. Our numerical experiment is devoted to answering a very specific question on the interaction between the coronal heating and the cross-section of the transition region. It is complementary to self-consistent 3D MHD descriptions of coronal boxes \\citep{Bingert2013a}.}  In Sec.\\ref{sec:model} we describe our model, the related equations and numerical code, the initial and boundary conditions, and how we set up the simulation of interest. In Sec.\\ref{sec:simul} the simulation and its results are illustrated and discussed in Sec.\\ref{sec:discuss}. ",
        "conclusions": "{\\bf F_{c}} =  \\frac{F_{sat}}{F_{sat} + \\left| {\\bf F_{class}} \\right|} {\\bf F_{class}} \\end{equation}  \\begin{equation} \\label{TC_flux_class_summary} {\\bf F_{class}} = k_{||} {\\bf \\hat{ b}} \\left( {\\bf \\hat{ b}} \\cdot {\\bf \\nabla} T \\right) + k_{\\bot} \\left[  {\\bf \\nabla} T - {\\bf \\hat{ b}} \\left( {\\bf \\hat{ b}} \\cdot {\\bf \\nabla} T \\right) \\right] \\end{equation}  \\begin{equation} \\label{TC_flux_abs_class_summary} \\left| {\\bf F_{class}} \\right| = \\sqrt{\\left( {\\bf \\hat{ b}} \\cdot {\\bf \\nabla} T\\right)^{2}(k_{||}^{2} - k_{\\bot}^{2}) + k_{\\bot}^{2} {\\bf \\nabla} T^{2}} \\end{equation}  \\begin{equation} \\label{TC_flux_sat_summary} F_{sat} = 5\\phi \\rho c_{iso}^3 \\end{equation}  \\noindent are the total pressure, and total energy density (internal energy $\\epsilon$, kinetic energy, and magnetic energy) respectively, $t$ is the time, $\\rho = \\mu m_{H} n_{H}$ is the mass density, $\\mu = 1.265$ is the mean atomic mass (assuming metal abundance of solar values; \\citealt{Anders_Grevesse_1989}), $m_{H}$ is the mass of hydrogen atom, $n_{H}$ is the hydrogen number density, $ {\\bf u}$ is the plasma velocity, ${\\bf g}$ is the gravity acceleration vector for a curved loop, ${\\bf I}$ is the identity tensor, $T$ is the temperature, ${\\bf F_{c}}$ is the thermal conductive flux (see Eq.~\\ref{TC_flux_pes_summary}, ~ \\ref{TC_flux_class_summary}, ~ \\ref{TC_flux_abs_class_summary}, ~ \\ref{TC_flux_sat_summary}), the subscripts $||$ and $\\bot$ denote, respectively, the parallel and normal components to the magnetic field, $k_{||} = K_{||} T^{5/2}$ and $k_{\\bot} = K_{\\bot} \\rho^2 / (B^2 T^{1/2})$ are the thermal conduction coefficients along and across the field, $K_{||}$ and $K_{\\bot}$ are constants, $c_{iso}$ is the isothermal sound speed, $\\phi = 1$ is a free parameter, and $F_{sat}$ is the maximum flux magnitude in the direction of ${\\bf F_{c}}$. $\\Lambda\\left( T\\right)$ represents the optically thin radiative losses per unit emission measure derived from CHIANTI v.$~7.0$ database \\citep{Chianti_I,Chianti_VI,Reale2012a} assuming coronal metal abundances \\citep{Feldman_1992}, and $H \\left( {\\bf s}, t \\right)$ is a  function of space and time describing the phenomenological heating rate (see Sect.\\ref{sec:simul}). We use the ideal gas law, $P = (\\gamma -1) \\rho \\epsilon$. We assume negligible viscosity and  resistivity, \\rev{except for those intrinsic in the numerical scheme}.     The calculations are performed using the PLUTO code \\citep{Mignone_2007,Mignone_al_2011}, a modular, Godunov-type code for astrophysical plasmas. The code provides a multiphysics, algorithmic modular environment particularly oriented toward the treatment of astrophysical flows in the presence of discontinuities as in the case treated here. The code is designed to make efficient use of massive parallel computers using the message-passing interface (MPI) library for interprocessor communications. The MHD equations are solved using the MHD module available in PLUTO, configured to compute intercell fluxes with the Harten-Lax-Van Leer approximate Riemann solver, while second order in time is achieved using a Runge-Kutta scheme. A Van Leer limiter for the primitive variables is used. The evolution of the magnetic field is carried out adopting the constrained transport approach \\citep{Balsara_Spicer_1999} that maintains the solenoidal condition $({\\bf \\nabla} \\cdot {\\bf B} = 0)$ at machine accuracy.  PLUTO includes optically thin radiative losses in a fractional step formalism \\citep{Mignone_2007}, which preserves the 2nd time accuracy, since the advection and source steps are at least $2^{nd}$ order accurate; the radiative loss $\\Lambda$ values are computed at the temperature of interest using a table lookup/interpolation method. The thermal conduction is treated separately from advection terms through operator splitting. In particular we adopted the super-time-stepping technique \\citep{Alexiades_al_1996} which has been proved to be very effective to speed up explicit time-stepping schemes for parabolic problems. This approach is crucial when high values of plasma temperature are reached (as during flares), explicit scheme being subject to a rather restrictive stability condition (i.e. $\\Delta t \\le (\\Delta x )^{2} / 2 \\eta $ where $\\eta$ is the maximum diffusion coefficient), as the thermal conduction timescale $\\tau_{cond}$ is shorter than the dynamical one $\\tau_{dyn}$ (e.g. \\citealt{Orlando_al_2005_A,Orlando_al_2008}).     \\subsection{The loop setup} \\label{sec:setup}    We address a typical active region loop, with half length $L = 3 \\times 10^{9}$ ~ cm and  temperature  $T \\sim 3 \\times 10^{6}$ ~ K.  We choose as initial plasma condition a plane-parallel loop atmosphere with the temperature and density profile from a model of hydrostatic loop \\citep{Serio_1981_bis} with an apex temperature about $8 \\times 10^{5}~K$. For the chromospheric part of the loop we use a hydrostatic atmosphere with a uniform temperature (i.e. $10^{4}$ K).  The temperature spans from $10^{4}$ K in the chromosphere to $8\\times 10^{5}$ K in the corona. The density spans from $\\sim 10^{14}$~cm$^{-3}$ in the chromosphere to $\\sim 10^{8}$~cm$^{-3}$ in the corona. This atmosphere is much more tenuous and cool than the subsequently heated one. The corona and the chromosphere are connected through a thin transition region where the temperature jumps from $10^4$ K to $10^6$ K in less than $10^8$ cm.  The computational domain (see Fig.$~\\ref{Fig:resol_init}$) extends from  $-3.1 \\times 10^{9} ~ cm$ to $3.1 \\times 10^{9} ~ cm$ in the $z$-direction (i.e along the loop axis), and from  $r=0$ to $r = 3.5 \\times 10^{9} ~ cm$ in the $r$-direction (i.e across the loop).  High resolution ($dr \\sim dz \\sim 3 \\times 10^{6}$ cm) is needed to describe appropriately the thin transition region from the chromosphere to the corona. The grid that we use is non-uniform but fixed. Fig.~\\ref{Fig:resol_init} shows the pixel size map in the domain. We have the maximum resolution around the transition region ($\\left| z \\right| \\approx 2.4 \\times 10^{9} ~ cm$, layers where the colors change abruptly in Fig.~\\ref{Fig:resol_init}) in the vertical direction and along the loop axis in the $r$-direction ($r \\approx 0$).   \\begin{figure}[!ht]               % \\centering {\\includegraphics[width=10cm]{fig1.eps}} \\caption{\\small Top: Density (left, $[cm^{-3}]$) and temperature (right, $[K]$) color map (logarithmic scale) before the loop heating is switched on. The chromosphere is red in the density map and blue in the temperature map. The magnetic field lines are also marked (white lines). Bottom: pixels size map in the computational domain. Each box encloses 10$\\times$10 grid points. } \\label{Fig:resol_init} \\end{figure}     A fundamental ingredient of our study is that the loop where we inject the heating expands from the chromosphere to the corona. We want that the magnetic field has this expansion already at the beginning of the heating release. Therefore, our strategy involves two steps: first, we build the magnetic configuration in a tenuous, static and plane parallel atmosphere, i.e. we build an ``empty'' loop. Then, we inject the heating in a localized region of the domain to switch the loop on. For the first step, we start with a magnetic field everywhere in the vertical direction, i.e. linking the two dense independent chromospheres. {\\it The magnetic field and the plasma density are more intense around the central axis, i.e. around $r=0$}. This configuration is not at equilibrium, and a preliminary simulation is devoted to let the configuration relax to equilibrium. The central initial density is chosen so that at equilibrium it is approximately equal to that of the surroundings.  In the tenuous corona, around the central axis the magnetic field pressure is much higher than the plasma pressure, i.e. the ratio of thermal pressure and magnetic pressure $\\beta \\ll 1$. As a result, the magnetic field expands considerably in the corona. The magnetic pressure is instead only a small perturbation in the chromosphere ($\\beta \\gg 1$) and it is left almost unchanged there. We let the whole system evolve freely, keeping a small heating that compensates for radiative and conductive losses. A quasi-steady state is reached after $t \\sim 4000$ s (the Alfven and sound crossing time are $\\sim 50$ s and $\\sim 1000$ s, respectively). The result of this preliminary step is a steady and stable loop atmosphere, with the strong expansion of the magnetic field lines from the chromosphere to the corona around the loop central axis, as we supposed (Fig.~\\ref{Fig:resol_init}). The magnetic field intensity decreases from a few hundreds $G$ in the chromosphere to a few $G$ in the corona (sufficient to confine the loop plasma). As shown in Fig.~\\ref{Fig:resol_init}, the area expansion is a factor of 6 from the bottom of the chromosphere to the top of the transition region, another factor 2 in the first 3000 km above the transition region, and another factor 2 to the top of the loop (middle of the domain).  The plasma atmosphere also readjusts to a new equilibrium: in Fig.~\\ref{Fig:resol_init} we see that in the magnetically-expanded configuration the atmosphere is no longer strictly plane parallel both in the temperature and in the density, but the inhomogeneities are small and leave the corona everywhere very tenuous and quite cooler than 1 MK. We also see some ripples in the chromosphere that do not affect the evolution when the loop heating is on. All plasma motions are a few km/s at most.  Many different initial conditions have been tested, finding that the resolution in the transition region is critical for accuracy.  We adopt axisymmetric boundary conditions at $r = 0$, i.e. along the symmetry axis, reflective boundary conditions at $r = 3.5 \\times 10^9$ cm and reflective boundary conditions but with reverse sign for the magnetic field components at $z = \\pm 3.1 \\times 10^9$ cm.     \\label{sec:discuss}  We investigated an MHD model of an active region loop, and in particular the variation of the loop cross-section in the moss region.  To study the low-lying parts of the loops in high detail, the assumption of constant cross-section all along the loop cannot hold, because it is known that the loop expands considerably going up from the transition region to the corona (e.g., \\citealt{Gabriel_1976}). We could not help including this effect in the modeling, and how it changes in time and with the plasma $\\beta$,  and this required a proper time-dependent magnetohydrodynamic description.  Our model considers the time-dependent MHD equations in a 2D cylindrical geometry including the gravity (for a curved loop), the thermal conduction, the coronal heating, and the radiative losses from optically thin plasma.  We first created the proper loop topology, with an expanding cross-section from the chromosphere to the corona, ready for the injection of heating. Then,  we heated our loop with a large-scale slowly-changing (quasi-steady) heating (e.g.,  \\citealt{Winebarger_al_2011,Warren_al_2010_a,Warren_al_2010_b}). We focussed on the structure of the so-called moss in active regions, i.e. the thin loop layers at temperature ranging between about 0.1 and 1 MK, that are generally interpreted as the footpoints of hotter loops. We measured the position and size of the moss layers.  From Fig.\\ref{Fig:moss} we have seen that the position and size of the moss region change with time. Consequently also the emission from moss varies.  The primary result from our simulations is that the area of the transition region can change substantially if there is a sizable change in the quasi-steady heating rate. For example, the area at 0.5 MK changes by about 40\\% as the peak temperature in the flux tube varies between 1 and 4 MK. The plasma beta at the relevant heights is low enough that the magnetic field is minimally influenced by the enhanced gas pressure. The dominant effect of the pressure is to force the transition region downward to a place where the tube is more highly constricted. This alters the DEM($T$) curve, since the shape of the curve, including the ratio of coronal to transition region emission, depends on how the flux tube area varies with temperature. It may not be reasonable to independently specify a particular $A(T)$ dependence when modeling the solar atmosphere. \\cite{Warren_al_2010_b} found an $A(T)$ that reproduces the observed DEM($T$) of an active region core, but it is not obvious that this A(T) is consistent with the assumed steady heating rate.  \\rev{Our model is simplified in several respects. Our 2D cylindrical description assumes a cylindrical symmetry around the loop central axis. Although there are some indications that this might not always be the case \\cite{Malanushenko2013a}, we do not expect considerable differences of the results for moderate deviations of the shape of the cross-section. The structure and energy balance of the chromosphere are also simplified into an isothermal structure with no radiation transport, but our study focusses on layers at much higher temperature that should not be highly affected by this approximation. Finally, although our model implies a full MHD description, still it is not entirely self-consistent, since the heating is imposed artificially. A fully consistent model is out of the scope of this work, since we focus on the very specific question of the reaction of the low corona to heating.}   \\begin{figure}[!ht]               % \\centering \\centering \\subfigure[] {\\includegraphics[width=7.5cm]{fig5a.eps}} \\subfigure[] {\\includegraphics[width=7cm]{fig5b.eps}} \\caption{\\small \\revb{(a) Map of synthetic emission in the Fe IX 171 \\AA\\ line at t = 0, in the same enlargement as in Fig. 4. The emission is normalized to its maximum and the color scale is shown in panel (b).} (b) Time(x-axis)-space($r$, y-axis) plot of the synthetic emission in the Fe IX 171 \\AA\\ line. We plot the radial profile of the emission in the vertical direction at any time. Each profile is normalized to its maximum. } \\label{Fig:fe9} \\end{figure}  Another interesting result is the magnitude of the flux tube area expansion between the transition region and the corona. The area increases by approximately a factor of 2 to 3 over the first 3000 km above the chromosphere in our model, depending on the peak temperature. This is much less than in the models of \\cite{Gabriel_1976,Athay1981a,Dowdy1986a,Warren_al_2010_b}. Those models have a \"throat\" of rather extreme expansion, which we suggest may not be realistic. The Gabriel and Athay models use a potential field and assume that the flux tube is highly constricted at the {\\it top} of the chromosphere. They further assume that the flux tube is isolated, with no surrounding field, which is probably a poor assumption, as discussed in \\cite{Dahlburg_2005}. Our \\rev{\"$\\beta$-sensitive\"} MHD model suggests that throats of extreme expansion are unlikely, although convective motions, not included here, may serve to enhance the effect. We note that the dependence of $A(T)$ on heating rate would be even stronger if such throats did exist.  \\bigskip"
    },
    "1402/1402.0612_arXiv.txt": {
        "abstract": "Faraday tomography is a powerful method to diagnose polarizations and Faraday rotations along the line of sight. The quality of Faraday tomography is, however, limited by several conditions. Recently, it is reported that Faraday tomography indicates false signals in some specific situations. In this paper, we systematically investigate the condition of the appearance of false signals in Faraday tomography. We study this by pseudo-observing two sources within a beam, and change in the intrinsic polarization angles, rotation measures, intensities, and frequency coverage. We find that false signals arise when rotation measure between the sources is less than 1.5 times the full width at half maximum of the rotation measure spread function. False signals also depend on the intensity ratio between the sources and are reduced for large ratio. On the other hand, the appearance of false signals does not depend on frequency coverage, meaning that the uncertainty should be correctly understood and taken into consideration even with future wide-band observations such as Square Kilometer Array (SKA). ",
        "introduction": "Cosmic magnetic fields are interesting topics for modern astrophysics and cosmology, because magnetic fields could affect the formation of cosmic structures and imprint themselves on histories and properties of cosmic evolutions. For example, magnetic fields are essential for instabilities in galactic gaseous disks \\citep[e.g.,][]{mac13}, play crucial roles in structures, particle accelerations, and radio emissions in galaxy clusters \\citep[e.g.,][]{tak08,fuj13}, record evolutions of turbulence in the cosmic web \\citep[e.g.,][]{ar10,ar11}, and affect cosmic ray or $\\gamma$-ray propagation such as pair echo \\citep{tkm09,ner10,tak12,tak13}.  Faraday rotation, the rotation of the polarization angle of electromagnetic waves traveling in magnetized plasma, is a standard tool to measure Galactic and extragalactic magnetic fields \\citep{sch09,wol10,nou12,bre11,gae05,bec09,fle11}. For a line of sight (LOS) toward a single polarized source, Faraday rotation can be written as \\begin{equation}\\label{eq1} \\chi(\\lambda^2) = \\chi_0 + \\phi \\lambda^2, \\end{equation} where $\\lambda$ is the wavelength, $\\chi$ and $\\chi_0$ are the rotation angles at the observer and the source, respectively, and $\\phi$ is the Faraday depth or Faraday rotation measure (RM) which is an integration of the LOS component of magnetic fields from the source to the observer with a weight of electron density. Notice that \"Faraday depth\" does not correlate with physical distance of source and sign which shows the direction of magnetic field toward observer dramatically changes the behavior of polarization angle and Faraday dispersion function (see Section 2).   Faraday depth can be estimated based on the linear relation between $\\lambda^2$ and $\\chi$ in equation (\\ref{eq1}). The estimation, however, becomes uncertain if there exists multiple sources within a beam, in the case $\\lambda^2$ and $\\chi$ does not display a linear relation \\citep[see e.g.,][]{bd05}. This problem is generally inevitable in studies of extragalactic magnetic fields, since we observe extragalactic sources viewed through Galactic emissions \\citep{bre11}. Therefore, we need more sophisticated methods to diagnose polarized emissions and Faraday depths along the LOS.  One of the sophisticated methods is so-called RM synthesis or Faraday tomography proposed in radio astronomy \\citep{bu66,bd05,and12,bec12,bel12,and13}, which has been applied by many authors \\citep{hea09,sch09,wol10,fri11,iac13}. This method transforms polarization spectra as a function of wavelength into those as a function of Faraday depth. For example, \\cite{osu12} analyzed point sources with Faraday tomography and spectral fitting with Faraday rotation models. They found that some of their polarization spectra could not be fitted with a single-source model but are well fitted with a multiple-source model. This means that they resolved the point sources by their polarization spectra. This demonstrates that Faraday tomography is effective for cases with multiple sources within a single beam.  Since the quality of Faraday tomography primarily depends on frequency coverage of the data, Faraday tomography is expected to become a major tool of radio polarimetry using future facilities such as Square Kilometer Array (SKA). Recently, however, \\cite{far11} reported some ambiguities associated with Faraday tomography. They considered two sources located very closely each other in Faraday depth, and varied the difference of intrinsic polarization angles of the two sources. They found that Faraday tomography could resolve two components correctly in some cases but, in other cases, it erroneously detected three components including false signals. This phenomenon, called RM ambiguity, would substantially degrade Faraday tomography unless we understand it.  In this paper, we extend the previous work and examine the conditions of the appearance of false signals more systematically. We investigate the dependence of false signals on intrinsic polarization angles, rotation measures, and intensities of the two sources. We also change frequency coverage based on planned observations with Australian SKA Pathfinder (ASKAP) and SKA. In section 2, we describe our model and calculations. We study the conditions of false signals in section 3, and give a summary and discussion in section 4. ",
        "conclusions": "Faraday tomography combined with RM CLEAN is a powerful method to investigate the distribution of polarized sources along the LOS. However, in the presence of multiple sources, large false signals can appear in the middle of the correct signals in $\\phi$ space. In this paper, we studied the condition of the appearance of false signals, focusing on the difference in the intrinsic polarization angles, the separation in $\\phi$ space and the intensity ratio. We simulated the observations of ASKAP and SKA assuming two polarized sources and found that false signals can appear for a separation as large as 1.5 FWHM of the RMSF and the false signals exceed the correct signals for a separation smaller than 1.1 FWHM. The intensity ratio is also a important factor and large ratio tends to reduce the RM ambiguity. It should be noted that these results does not depend on the bandwidth of the telescope so that even future wide-band observation would suffer from RM ambiguity.  It should be noted that an absolute value of $\\Delta\\phi$ for which we see false signals is decreased as decreasing the FWHM. More specifically, false signals arise for the relative RM smaller than 22.26 rad/m$^2$ with ASKAP, but than 0.189 rad/m$^2$ with SKA. Therefore, false signals would become a minor issue in the SKA era while we study RMs of a few rad/m$^2$ or larger.  RM CLEAN has three parameters, the gain factor $\\gamma$, threshold intensity, and maximum iteration. We also have studied the cases with different parameters, and found that larger gain factor tends to slightly increase the range of $\\Delta\\chi_0$ where false signals arise. On the other hand, threshold intensity and maximum iteration in this work are rather conservative values and the results were not changed significantly. Note that threshold intensity is usually referred to RMS noise level of the observations. Therefore, if we study polarization sources with less than mJy, the threshold intensity in this work could reach the RMS noise levels of ASKAP and SKA, and thus should increase the threshold. For such cases, we may need more careful arguments for the RM CLEAN parameters.  Finally, Stokes' QU-fitting is another powerful method to probe cosmic magnetism with polarization observation \\citep{osu12,ide13}. In this method, fitting of polarization spectrum is performed assuming a model of source distribution with free parameters. Combination of QU-fitting and Faraday tomography would be effective to identify and avoid RM ambiguity, and detect correct polarized sources. This possibility will be pursued elsewhere in near future."
    },
    "1402/1402.7174_arXiv.txt": {
        "abstract": "We present the results of high resolution VLBI observations at 1.6 and 4.9~GHz of the radio-loud Seyfert galaxy, Mrk\\,6. These observations are able to detect a compact radio core in this galaxy for the first time. The core has an inverted spectral index ($\\alpha^{1.6}_{4.9}$=+1.0$\\pm$0.2) and a brightness temperature of $1\\times10^8$\\,K. Three distinct radio components which resemble jet elements and/or hot spots, are also detected. The position angles of these elongated jet elements point, not only to a curved jet in Mrk\\,6, but also towards a connection between the AGN and the kpc-scale radio lobes/bubbles in this galaxy. Firmer constraints on the star formation rate {provided by} new {\\it Herschel} observations (SFR $<0.8$~M$_{\\sun}$~yr$^{-1}$) make the starburst-wind powered bubble scenario implausible. From plasma speeds obtained via prior {\\it Chandra} X-ray observations, and ram pressure balance arguments for the ISM and radio bubbles, the north-south bubbles are expected to take {$7.5\\times10^6$~yr} to form, and the east-west bubbles $1.4\\times10^6$ yr. We suggest that the jet axis has changed at least once in Mrk\\,6 within the last {$\\approx10^7$~yr}. A comparison of the nuclear radio-loudness of Mrk\\,6 and a small sample of Seyfert galaxies with a subset of low-luminosity FRI radio galaxies reveals a continuum in radio properties. ",
        "introduction": "{Markarian\\,6 is a nearby\\footnote{$z$ = 0.018813, luminosity distance = 80.7 Mpc} bulge-dominated lenticular galaxy hosting a Seyfert 1.5 type active galactic nucleus (AGN)}. Lenticular galaxies are considered to be intermediate between spiral and elliptical galaxies; they have a disk but have either consumed or lost most of their interstellar medium (ISM), so that there is little ongoing star formation activity. {At radio frequencies,} Mrk\\,6 exhibits a highly complex morphology: two pairs of self-similar bipolar bubble-like radio structures {that are} oriented nearly perpendicular to each other {are clearly visible} \\citep{Kharb06}. While starburst superwinds have been suggested to be {mainly responsible for} the creation of lobe/bubble-like structures in Seyfert galaxies \\citep[e.g.,][]{Baum93}, several pieces of evidence point towards an AGN {origin} \\citep[e.g.,][]{Colbert96,Elmouttie98,Gallimore06,HotaSaikia06,Kharb06,Kharb10a}.  The complex radio morphology of Mrk\\,6, along with a high degree of polarization ($\\sim50$\\% at the edges) and an AGN jet-like spectrum {($\\alpha\\sim-0.8$; defined throughout as $S_\\nu\\propto \\nu^{\\alpha}$),} led \\citet{Kharb06} to suggest that the edge-brightened bubbles were the projections of precessing radio jets. The precessing jets had to shut-off and restart in a nearly perpendicular direction to account for both sets of bubbles. The dominant role played by the AGN in the creation of the radio bubbles was consistent with the low star-formation rate (SFR) of $\\le5.5$~M$_\\sun$~yr$^{-1}$ derived from the {\\it Spitzer} Space Telescope data of \\citet{Buchanan06}: an SFR of $\\approx33$~M$_\\sun$~yr$^{-1}$ was required for the bubbles to have been inflated by supernova explosions. New {\\it Herschel} telescope data place even more stringent constraints on the SFR, as discussed ahead.  High-resolution radio observations with the MERLIN\\footnote{Multi-Element Radio-Linked Interferometer Network} interferometer were able to identify individual components in the north-south kpc radio jet \\citep{Kukula96}. As the MERLIN observations were carried out at a single frequency, \\citet{Kukula96} could not identify a flat spectrum radio core in Mrk\\,6. On the basis of the similar position angles (PAs) of the jet and the brightest regions along the edges of the large north-south lobes, they suggested that these bright regions were hot spots created by the north-south jet, where most of the bulk energy of the jets was dissipated; the jet plasma then drifted along the local host galaxy density gradient along the minor axis, forming the rest of the radio lobes. This of course cannot explain the east-west bubbles in this source, because they lie close to the major axis of the galactic bulge {(and the galactic disk)}, where the density gradient is not so pronounced. However, since Kukula et al. did not actually detect the full extent of the east-west bubbles in their observations, this question did not arise. The precessing jet scenario invoked by \\citet{Kharb06} suggested on the other hand, that the jet did not dissipate all its energy in these hot spots but rather continued to advance beyond them to create the 8-shaped radio structures. Furthermore, the precession model indicated that the north-south kpc jet was more closely associated with the east-west bubbles, rather than the north-south bubbles \\citep[see Figure~6,][]{Kharb06}.  In order to re-examine these issues, as well as to search for a flat or inverted spectrum radio core, we observed Mrk\\,6 with the Very Large Baseline Array (VLBA) at 1.6 and 4.9~GHz. We present here the results from these observations. We also discuss results from two epochs of archival 8.4~GHz VLBA data on Mrk\\,6. We have used radio bubbles and lobes interchangeably in the text. However, the lobes in Seyfert galaxies differ from those found in powerful {Fanaroff-Riley type II radio galaxies \\citep[FRII;][]{Fanaroff74}} in that they are edge-brightened on all sides \\citep[e.g.,][]{Mingo12} and not just at the leading edges like in FRIIs \\citep[e.g.,][]{Kharb08}.  At the distance of Mrk\\,6, $1\\arcsec$ corresponds to 377 pc, for $H_0$ = 71 km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{mat}$ = 0.27, $\\Omega_{vac}$ = 0.73. ",
        "conclusions": "In \\citet{Kharb06}, we had estimated that the star formation rate derived from {\\it Spitzer} observations was far below that required to inflate the radio bubbles. Additional evidence for a comparatively low SFR comes from new {\\it Herschel} data which we obtained from the {\\it Herschel} archive and reduced using the {\\it Herschel} analysis software HIPE (see Section~2 for details). We present in Figure~\\ref{figsed} the broad-band SED of Mrk\\,6 with new {\\it Herschel} infrared data using the PACS and SPIRE instruments \\citep{Poglitsch10,Griffin10}. {No previous data existed in this part of the spectrum and this figure clearly shows the importance of the new {\\it Herschel} data for tracing the infrared peak of the SED, which originates in AGN and/or star formation heated dust emitting in the infrared.} These {\\it Herschel} data have the advantage over previous {\\it Spitzer} measurements in that the star formation is thought to be the dominant heating mechanism of the dust that produces emission at these longer wavelength in the far-infrared. At {\\it Spitzer} wavelengths (i.e. 24$\\mu$m) a considerable contribution to the 24$\\mu$m luminosity may be expected from the AGN instead. Therefore, the SPIRE 250$\\mu$m band can give a better estimate of the SFR than was previously possible with {\\it Spitzer}. Calculating {the luminosity at 250$\\mu$m (L250$\\mu$m)} for Mrk\\,6 and using the empirical relation derived by \\citet{Hardcastle13} for L250$\\mu$m versus SFR, measured from optical spectral modelling, we find a low SFR of 0.8~M$_\\sun$ yr$^{-1}$. This SFR is also an upper limit because the 250$\\mu$m flux can originate both from cool dust heated by the old stellar population as well as from young star forming regions. It is also noteworthy that Mrk\\,6 has no strong evidence for star formation tracing polycyclic aromatic hydrocarbon (PAH) features in contrast to the majority of Seyfert galaxies \\citep{Gallimore10}. Therefore, for the remaining discussion we have assumed that AGN jets are primarily responsible for the creation of the radio bubbles. We discuss the morphology of the radio jets, their connection to the kpc-scale bubbles, and the bubble ages in the following sections. Finally we discuss the radio-loudness of Mrk\\,6 and other Seyfert galaxies and compare them to those of low luminosity FRI radio galaxies.  \\subsection{Changes in Jet Direction: Interaction with the Medium} \\citet{Kinney00} used the distribution of radio jet orientation in Seyfert galaxies to show that the accretion disks are commonly misaligned with the host galaxy disk. There are several reasons why this might be the case, including misaligned inflow of gas from a minor merger. The misaligned gas inflow will change the accretion disk axis, and the blackhole will re-align with the disk via Lense-Thirring precession \\citep{Bardeen75,Rees78}. Minor mergers could be responsible for the creation of randomly-oriented short-lived accretion disks in Seyfert galaxies, that could in turn lead to jet precession \\citep[e.g.,][]{Volonteri07}. Another mechanism for a change in the jet axis would be a merger with another black hole \\citep[e.g.,][]{Merritt02}. {For the brightest cluster galaxy of RBS797, which exhibits orthogonal radio outflows, a jet flip due to the presence of a binary black hole has been suggested by \\citet{Gitti13}.} ISM rotation and ram pressure bending have also been invoked to explain the curved radio jets in some Seyfert galaxies \\citep[e.g.,][]{Wilson82}. While jet flips alone may arise from a minor merger event, both jet precession and jet flips, as discussed for Mrk\\,6 ahead, can naturally occur due to the presence of binary black holes.  The archival {\\it Hubble Space Telescope (HST)} Advanced Camera for Surveys (ACS) F330W image\\footnote{Program ID 9379, PI - Henrique Schmitt, http://archive.stsci.edu/cgi-bin/mastpreview?mission=hst$\\&$dataid=J8EK17011} of Mrk\\,6 shows the presence of a tail-like feature towards the north-west, which could be a signature of past merger activity in Mrk\\,6. The emission line images of [OIII] and [OII] from the Faint Object Camera (FOC) aboard the {\\it HST} show that there is dense emission line gas in the direction of the nucleus in Mrk\\,6, which could provide the clouds which merge with the accretion disk. The emission line images show a clear correspondence with the southern part of the MERLIN radio jet \\citep[see][]{Capetti95}. Capetti et al. however used a weak radio component in between S1 and S2 as the radio core to align their emission-line image. {When the detected radio core is properly aligned, we find that the radio component S2 lies at the outer edge of the bend observed in the S-shaped emission line region (see Figure~\\ref{figfoc}).} This close correspondence could be indicative of interaction between the radio ejecta and the surrounding medium \\citep[e.g.,][]{Capetti95}. {\\it Chandra} and {\\it XMM-Newton} observations of Mrk\\,6 have shown that the X-ray emitting gas has a complex structure \\citep{Schurch06,Mingo11}. This is also consistent with a jet that is changing its direction of propagation, interacting with the surrounding ISM.  \\subsection{A Curved Jet due to Precession ?} \\label{secprec} We first examine the precession model to explain the S-shaped jets in Mrk\\,6. The precession model is in keeping with the basic idea proposed by \\citet{Kukula96}, namely, that the north-south jets are producing the hot spots along the edges of the kpc-scale north-south bubbles. We note that Kelvin-Helmholtz instabilities that arise in relativistic jets due to some perturbation at the jet nozzles could also produce helical twisting patterns in jets \\citep[e.g.,][]{Hardee87}. At the end of the section, we examine an alternate model proposed by \\citet{Mingo11} to explain the Mrk\\,6 lobes.  Following the relations of \\citet{Hjellming81}, we created a precessing jet model, and superimposed it on the VLBA and VLA radio images of Mrk\\,6 (see Figure~\\ref{figprec}). {The 20cm VLA image is reproduced from \\citet{Kharb06}.} The new core position was chosen as the origin for this jet model. The projected precessing jet model matches the radio features well. The best fit model parameters {for a clockwise rotating jet} are: jet speed = 0.13$c$ ($\\sim3.9\\times10^4$~km~s$^{-1}$), jet PA = $-87\\degr$, jet inclination = 144$\\degr$, jet precession cone half-opening angle = 37$\\degr$, and jet angular velocity = 1.5$\\times10^{-6}$~rad~yr$^{-1}$. It is important to note that the above set of parameters is not unique, and another combination of the parameter set could also result in a good fit to the apparent radio morphology. This is also demonstrated by the fact that the best fit parameters presented here differ from those derived in \\citet{Kharb06}, although different spatial scales were examined in the two cases. One therefore needs independent measurements of as many parameters as possible (e.g., jet speed, jet inclination), to realistically constrain the precession model. However, it is important to note that precessing jets have independently been inferred in Seyfert galaxies, through core flux and PA variability studies, carried out by multi-epoch VLBI obervations \\citep[e.g., III\\,Zw\\,2 \\& M\\,82;][]{Brunthaler05,MartiVidal11}.  For the best fit jet speed and inclination obtained here, a jet-to-counterjet intensity ratio ($\\equiv[1+\\beta\\cos\\theta/1-\\beta\\cos\\theta]^p$) {of 0.6 is expected for a jet structural parameter, $p$ of 2.3, and 0.5 for a $p$ of 3.3 \\citep[$p=2-\\alpha$, or $3-\\alpha$ for a continuous or ``blobby'' jet, respectively;][]{UrryPadovani95}.} It is interesting though that no jet or counterjet emission is detected to within 22 parsec of the core. While such a gap is traditionally considered to be a signature of relativistic deboosting in fast jets that are lying close to the plane of the sky, it is difficult to assume the same for Seyfert galaxies, since Seyfert jets are expected to be only mildly relativistic \\citep[e.g.,][]{Bicknell98,Middelberg04}.  {The way it is} depicted in Figure~\\ref{figprec}, the jet has been ``on'' for 1.4$\\times10^5$ yr. This time constraint follows from the basic idea proposed by \\citet{Kukula96} that the north-south jet gets disrupted after forming hot spots along the edges of the north-south lobes. Therefore, 1.4$\\times10^5$ yr was the time the jet must have been ``on\" in the precession model, in order for it to be disrupted around the position of the bright edges or hot spots of the north-south bubbles {(as determined visually from Figure~\\ref{figprec})}. From this best fit model, it {is apparent} that component N is not a part of the jet that connects the parsec scale emission to the kpc-scale north-south bubbles, but is rather a part of the east-west bubbles. This is also seen by the PA of its extension in the left panel in Figure~\\ref{figwhole}. The east-west extensions in Figure~3 of \\citet{Kukula96}, which are clearly part of the east-west bubbles, also have the same PA. The fact that this radio component has the steepest spectra (actually only a lower limit) of the three detected components, could be consistent with it arising from a different AGN activity episode, compared to components S1 and S2. The steep spectrum of component N could suggest electron radiative decay in older radio emission. However, this is inadequate evidence by itself, since the spectrum of component S2 is also steep. {While the MERLIN observations of Kukula et al., were unable to resolve components S1 and S2 for the purpose of estimating the 1.6$-$4.9~GHz spectral indices, they derived an average $\\alpha=-0.8$ for that part of the jet. Their component 6 (component N here) was easily resolved and had a somewhat steeper spectral index of $-0.9$.}  Apart from the redundancies present in the best-fitting model, another caveat of the precession model, as presented here, is that it suggests that the jet is currently powering the larger north-south lobes, rather than the smaller but brighter east-west lobes. While the extended narrow line region, which is likely being photoionized and/or shock-ionized by the AGN, does lie along the north-south direction, consistent with this being the current jet direction; the size and surface brightness of the east-west lobes are suggestive of them being the {ones that are newly formed.}  Another model, suggested by \\citet{Mingo11}, proposes that the bubbles are edge-brightened lobes created by the AGN that are expanding into the ISM. This scenario also requires a flip in the AGN {jet} axis. This model is consistent with the X-ray observations which detect shocked gas (albeit only in small regions) ahead of the radio lobes. A drawback of this model is that it cannot explain why the upper edge of the northern lobe and the lower edge of the southern bubble is relatively brighter than the opposite edge. This effect is much more pronounced in the east-west bubbles where the lower edge of the east bubble and upper edge of the west bubble is brighter. If we assume that the ram pressure of the lobe against a counter-clockwise rotating ISM is responsible for the selective edge brightening in the north-south lobes, then the ISM rotation would need to be clockwise to explain the bright opposite edges in the east-west lobes. This could be a big drawback for the edge-brightened lobe model. Selective edge brightening can more easily be explained as directional, {albeit} curved, AGN outflows.  \\subsection{Constraints on Bubble Ages from Radio and X-ray Observations} \\citet{Mingo11} estimated that the total work available from the plasma filling the north-south lobes was $\\approx4\\times10^{56}$~erg. For the AGN jet with a total radio luminosity of say, $\\approx10^{40}$~erg~s$^{-1}$ {\\citep[][Table~3]{Kharb06},} and kinetic luminosity of $10^{42}$~erg~s$^{-1}$ \\citep[assuming an efficiency of conversion of 1 percent;][]{Odea85,Birzan07}, the jet would require $1.3\\times10^{7}$ years to power the north-south lobes. This is about 90 times longer than the time the jet was ``on'' in the precessing jet model. The two timescales would be comparable if the conversion efficiency was lower than 0.1 percent.  Based on the Mach number $M$ of the shock producing the X-ray emission at the outer edges of the north-south bubbles in Mrk\\,6 ($M=3.9$), \\citet{Mingo11} estimated that the speed of the plasma in the bubbles was around 1030$^{+75}_{-100}$ km~s$^{-1}$. From the {\\it Chandra} observations, they also estimated that the electron density of the ISM outside the bubbles varied from 8 - 40$\\times10^{-4}$~cm$^{-3}$. This implies that the ISM gas mass density {(obtained by multiplying the ISM electron density with the mass of the hydrogen atom = $1.67\\times10^{-24}$ g)} varies from 1.3$-$6.7$\\times10^{-27}$~g~cm$^{-3}$. If the minimum pressure of the north bubble varies from $1.2\\times10^{-12}$~dyne~cm$^{-2}$ to $6.2\\times10^{-11}$~dyne~cm$^{-2}$ for a filling factor of 1.0 and $10^{-3}$, respectively \\citep[see Table 3,][]{Kharb06}, then the velocity of the lobe plasma can be estimated assuming that the lobe is in ram pressure balance with the surrounding ISM. Using the relation $P_{min} = \\rho_{ISM} v_{lobe}^2$ and the ISM gas density from above, we find that $v_{lobe}$ varies between $\\sim$130 to 960 km~s$^{-1}$ for a gas density of 6.7$\\times10^{-27}$~g~cm$^{-3}$, and between $\\sim$300 to 2150 km~s$^{-1}$ for a gas density of 1.3$\\times10^{-27}$~g~cm$^{-3}$. Assuming that the plasma velocity inside the lobes was around 1000~km~s$^{-1}$, the north bubble of extent 7.7 kpc would take 7.5$\\times10^6$ years to form, while the east bubble of extent 1.4 kpc would take 1.4$\\times10^6$~yr to form for the same plasma speed. However, as {\\it Chandra} could not clearly resolve the X-ray emission from the east-west bubbles, the plasma speed and therefore the age is uncertain for the east-west structure. It is only clear that the jet axis in Mrk\\,6 has changed at least once within the last $\\approx10^7$~yr.  \\subsection{The Radio Loudness of Seyferts} Seyfert galaxies have been historically identified as radio-quiet AGN. \\citet{Kellermann89} defined the radio-loudness parameter, $R$, of an AGN to be the ratio of its radio flux density at 5~GHz to its $B$-band optical flux density at 4400 \\AA, with the radio-quiet/radio-loud AGN division set at $R\\sim10$. {As Seyferts reside in spiral or S0-type host galaxies, and radio-loud AGN are typically hosted by elliptical galaxies, explanations for the ``lack\" of large-scale radio emission in Seyferts, have focused on differences in their central engines, which are in turn dependent on their host galaxy types. {For} instance, differences in black hole masses and spins, both of which are different in spiral and elliptical galaxies, have been suggested to explain the differences in the radio-loudness properties \\citep[e.g.,][]{Sikora07}.} It is in this context that it is interesting to note that Mrk\\,6 has an $R$ value of $\\approx23$, making it a radio-loud Seyfert galaxy (the total 5 GHz flux density as measured by the single dish Green Bank Telescope for Mrk\\,6 is 115$\\pm$9 mJy \\citep{Gregory91}, while the $B$-band apparent magnitude is 14.88$\\pm$0.16 mag (NED)\\footnote{NASA/IPAC Extragalactic Database}). The mere existence of radio-loud Seyferts can provide strong constraints on the jet production mechanisms in AGN.  The inherent assumption in the estimation of $R$ above, is that the $B$-band luminosity is dominated by the AGN light, which is unlikely to be true for Seyfert galaxies. \\citet{HoPeng01} extracted the optical nuclear luminosities for a large number of Seyfert galaxies through a proper modelling of their galactic bulge emission in {\\it HST} images. Remarkably, after taking into account their optical nuclear luminosities, Ho \\& Peng found that a majority of Seyfert 1 galaxies were radio-loud. A similar conclusion was reached by \\citet{Terashima03} on the basis of the radio to X-ray flux density ratio for a sample of low luminosity AGN including Seyferts. We present in the Appendix a comparison of the radio-loudness parameters of Seyfert (1 and 2) galaxies belonging to the Extended 12$\\mu$m sample of \\citet{Rush93}, along with a sample of UGC FRI radio galaxies, having taken into account their optical nuclear luminosities. We estimate that, even without {acknowledging} the AGN-only contribution to the optical luminosity, $\\approx$11 percent of Seyfert galaxies (both type 1s and 2s) from the Extended 12$\\mu$m complete sample of \\citet{Rush93} can be classified as radio-loud following the Kellermann et al. definition. This fraction increases to $\\ge$70 percent, when the optical nuclear luminosities are considered (see Table~\\ref{tabradl}).  We have plotted in Figure~\\ref{figrl} the radio-loudness parameter values for these Seyfert galaxies, along with those of 21 UGC FRI radio galaxies \\citep{Verdoes99,Kharb12a}. There is an overlap between the $R$ values of this subset of Seyferts and those of low luminosity FRI radio galaxies. \\citet{Balmaverde06} and \\citet{Kharb12b} also found a continuum of properties in the host galaxy surface brightness profiles as observed by the {\\it HST}, which reflects on the galaxies' nuclear structures, for a small sample of Seyfert and FRI radio galaxies. Although on average Seyfert galaxies have less massive blackholes than FRI radio galaxies, there is an overlap in the blackhole mass ranges for the two classes: $\\sim10^7-10^8$\\,M$_\\sun$ for Seyferts \\citep{Wandel99}, and $\\sim10^8-10^9$\\,M$_\\sun$ for FRI radio galaxies \\citep{Woo02}.  There is also a continuum in the radio morphological properties of FRI galaxies and radio-loud Seyfert galaxies, although more sensitive observations are required to map the diffuse lobe emission in the latter. In the absence of sensitive radio observations of Seyfert galaxies, we can restrict our comparison to individual sources, which prevents us from claiming a statistically significant overlap in radio morphology between Seyfert galaxies and low luminosity FRIs. The radio structure in the Seyfert galaxy, Circinus, with hot spot like features immersed in bubble-like lobes \\citep[see][]{Mingo12}, resembles the hot spot-lobe structure observed in the FRI radio galaxy, Hercules A \\citep{Gizani03,ODea13}. While the radio bubbles of Mrk\\,6 and other Seyfert galaxies (e.g., NGC\\,6764) are typically a few kpc to $\\sim$10 kpc across, and are confined to the bulge of the host galaxy, bipolar bubbles of extent $>$100 kpc have been observed in the radio galaxy, NGC\\,1167 \\citep{Shulevski12}. Interestingly, the FRI radio galaxy with the lowest nuclear radio-loudness parameter in the sample under study, UGC\\,03695 ($R$=3775), has a radio structure very similar to that seen in Seyfert galaxies (see Figure~\\ref{figrlfr1}). These data are from our VLA program AB920 (J. Noel-Storr et al., in preparation). The lower frequency image (1.4~GHz) shows the bipolar bubble-like lobes, while the higher frequency image (4.9 GHz) shows a jet-like (i.e., elongated) structure pointing toward the western edge of the north lobe. This not only suggests a close connection between the jet-like structure and the bubble-like lobes, but also the presence of a misaligned jet. Small-scale jets that are misaligned with the radio lobes are a common characteristic of Seyfert galaxies.  S0 galaxies resemble elliptical galaxies in terms of their bulge properties, but additionally have low density gas and dust-rich disks in their {centres}. It is {then} plausible that the radio jets in Seyfert galaxies, which likely arise from the inner accretion disks \\citep[e.g.,][]{Nixon13} are affected by the presence of these kpc-scale dusty disks, and precess. Any kind of warp that develops in the disk close to the jet launching sites, as a result of misalignment with the large scale galactic disks, may affect the jets so that they are no longer straight. Moreover, the distance from the cores where the jets ``flare'', or decollimate and widen, like in FRI radio galaxies \\citep[e.g.,][the flaring distance is typically 1 kpc in FRIs]{Laing02}, could be shorter for Seyfert galaxies. Subsonic jets moving in a confined medium where the density varies more slowly that 1/R$^2$, where R is the distance from the galaxy centre, could produce bubble-like radio structures \\citep{Falle91}. Therefore, we propose that a combination of an early-flaring, misaligned jet in a confined medium could produce Mrk\\,6-like radio lobes, {and similar lobes observed in other Seyfert galaxies}."
    },
    "1402/1402.1491_arXiv.txt": {
        "abstract": "We demonstrate that growth of stellar bars in spinning dark matter halos is heavily suppressed in the secular phase of evolution, using numerical simulations of isolated galaxies. In a representative set of models, we show that for values of the cosmological spin parameter $\\lambda\\gtorder 0.03$, bar growth (in strength and size) becomes increasingly quenched. Furthermore, slowdown of bar pattern speed weakens substantially with increasing $\\lambda$, until it ceases completely. The terminal structure of the bars is affected as well, including extent and shape of their boxy/peanut bulges. The essence of this effect lies in the modified angular momentum exchange between the disk and the halo facilitated by the bar. For the first time we have demonstrated that a dark matter halo can emit and not purely absorb angular momentum. Although the halo as a {\\em whole} is not found to emit, the net transfer of angular momentum from the disk to the halo is significantly reduced or completely eliminated. The paradigm shift implies that the accepted view that disks serve as sources of angular momentum and halos serve as sinks, must be revised. Halos with $\\lambda\\gtorder 0.03$ are expected to form a substantial fraction, based on lognormal distribution of $\\lambda$. Dependence of secular bar evolution on halo spin, therefore, implies profound corollaries for the cosmological evolution of galactic disks. ",
        "introduction": "\\label{sec:intro}  Redistribution of angular momentum in astrophysical systems is a major driver of their dynamical and secular evolution. Galactic bars facilitate this process  by means of gravitational torques, triggered internally (spontaneously) or externally (interactively). Important aspects of stellar bar evolution are still being debated --- their origin and evolutionary changes in morphology, growth and decay, are not entirely clear. Theoretical studies of angular momentum redistribution in disk-halo systems have been limited almost exclusively to {\\it nonrotating} halos, following pioneering works on linear perturbation theory by Lynden-Bell (1962), Lynden-Bell \\& Kalnajs (1972), Tremaine \\& Weinberg (1984) and Weinberg (1985), which underscored the dominant role of orbital resonances. Numerical simulations have confirmed the angular momentum flow away from disks embedded in axisymmetric (e.g., Sellwood 1980; Debattista \\& Sellwood 1998, 2000; Tremaine \\& Ostriker 1999; Villa-Vargas et al. 2009, 2010; review by Shlosman 2013) and triaxial (e.g., El-Zant \\& Shlosman 2002; El-Zant et al. 2003; Berentzen et al. 2006; Berentzen \\& Shlosman 2006; Heller et al. 2007; Machado \\& Athanassoula 2010; Athanassoula et al. 2013) halos. Resonances have been confirmed to account for the lion's share of angular momentum transfer (e.g., Athanassoula 2002, 2003; Martinez-Valpuesta et al. 2006; Weinberg \\& Katz 2007), In this paradigm, the halo serves as the pure sink and the disk as the net source of angular  momentum.  However, realistic cosmological halos are expected to possess a net angular momentum, acquired during the maximum expansion epoch (e.g., Hoyle 1949; White 1978) and possibly during the subsequent evolution (Barnes \\& Efstathiou 1987; but see Porciani et al. 2002). Simulations have quantified the distribution of spin values, $\\lambda\\equiv J_{\\rm h}/\\sqrt{2} M_{\\rm vir}R_{\\rm vir}v_{\\rm c}$, for cosmological dark matter (DM) halos to follow a lognormal distribution, where $J_{\\rm h}$ is the angular momentum, $M_{\\rm vir}$ and $R_{\\rm vir}$ --- the halo virial mass and radius, and $v_{\\rm c}$ --- the circular velocity at $R_{\\rm vir}$,  with the mean value $\\lambda = 0.035\\pm 0.005$ (e.g., Bullock et al. 2001). Spinning halos can increase the rate of the angular momentum absorption --- an issue brought up by Weinberg (1985) but never fully addressed since. Only recently has it been confirmed numerically that the bar instability timescale is indeed shortened for $\\lambda>0$ (Saha \\& Naab 2013). But these models had been terminated immediately after the bar instability had reached its peak, and hence avoided completely the secular stage of bar evolution.  The $\\lambda=0$ halos consist of two populations of DM particles, prograde and retrograde (with respect to disk spin). The amount of angular momentum in {\\it each} of these populations can vary from zero for nearly radial orbits, to a maximal one for nearly circular orbits. (Both extremes are mentioned for pedagogical reasons only.) These extremes in angular momentum correspond to extremes in velocity anisotropy.  Various degrees of velocity anisotropy in the halo lie in between, and represent a rich variety of dynamical models. Stellar bars mediate the angular momentum transfer in such disk-halo systems with a broad range of efficiencies. The current paradigm of stellar bar evolution assumes an idealized version of a nonrotating DM halo which cannot account for the whole bounty of associated processes. We address these issues in a subsequent paper (in preparation).  In this Letter we demonstrate for the first time that secular growth of galactic bars in spinning DM halos is damped more strongly with increasing $\\lambda$, and this effect is the result of a modified angular momentum transfer. Section~2 describes our numerical methods. Results are given in section~3. ",
        "conclusions": ""
    },
    "1402/1402.1172_arXiv.txt": {
        "abstract": "We use the public low redshift {\\it Zurich ENvironmetal Study (ZENS)} catalog to study the dependence of the quenched satellite fraction at $10^{10.0} M_\\odot \\rightarrow 10^{11.5} M_\\odot$, and of the morphological mix of these quenched satellites, on three different environmental parameters: group halo mass, halo-centric distance and large-scale structure over-density. At a given galaxy stellar mass, the fraction of quenched satellites is more or less independent of halo mass and the surrounding large-scale structure overdensity, but increases towards the centres of the halos, as found in previous studies.  The morphological mix of these quenched satellites is, however, strikingly constant with radial position in the halo, indicating that the well-known morphology-density relation results from the increasing fraction of quenched galaxies towards  the centres of halos.  If the radial variation in the quenched fraction reflects the action of two quenching processes, one related to mass and the other to environment, then the constancy with radius of the morphological outcome suggests that both  have the same effect on the morphologies of the galaxies. The quenched satellites have larger $B/T$ and smaller half-light radii than the star-forming satellites. The bulges in quenched satellites have very similar luminosities and surface brightness profiles, and any mass growth of the bulges associated with quenching cannot greatly change these quantities. The differences in the light-defined $B/T$ and in the galaxy half-light radii are mostly due to differences in the  disks, which have lower luminosities in the quenched galaxies.  The difference in galaxy half-light radii between quenched and star-forming satellites is however larger than can be explained by uniformly fading the disks following quenching, and the quenched disks have smaller scale lengths than in star-forming satellites. This can be explained either by a differential fading of the disks with galaxy radius or if disks were generally smaller in the past, both of which would be expected in an inside-out disk growth scenario.  The overall conclusion is that, at least at low redshifts, the structure of massive quenched satellites at these masses is produced by processes that operate before the quenching takes place.  A comparison of our results with semi-analytic models argues for a reduction in the efficiency of  group halos in quenching their disk satellites and for mechanisms to increase the $B/T$ of low mass quenched satellites. ",
        "introduction": "The star formation rate (SFR) and the morphology of a galaxy are key diagnostics of its evolutionary stage. These properties are broadly correlated, in the sense that  galaxies with dominant disks tend to have higher specific SFR (sSFR=SFR/$M_{galaxy}$, with $M_{galaxy}$ the galaxy stellar mass) than galaxies with a pronounced spheroidal component.  There is solid observational evidence that the cessation of star formation in some galaxies, which results in the emergence of  `quenched'  passive galaxies, depends on both galaxy mass and environment  \\citep[e.g.][]{Dressler:1980kx,Balogh:2004fk,Baldry:2006uq,Kimm:2009ij}. In the SDSS \\citep{York:2000fk} and zCOSMOS \\citep{Lilly:2007uq,Lilly:2009kx} samples, the effects of galaxy mass and environment in the quenching of galaxies appears to be separable, both in the local universe, and at least up to redshift $z=1$  (\\citealt{Peng:2010fk},  Kov\\^ac et al.\\ 2013). In other words it is possible to write the fraction of galaxies that survive as blue star-forming galaxies as the product of two functions, one a function of stellar mass only, and the other a function of  `environment'  only.  Each of these functions is therefore the extra effect of mass, or environment, over the effects of the other.   In these works, the environmental measure used is an Nth-nearest neighbour estimate of the local density of galaxies.  This separability of the quenching effects of galaxy mass and environment argues has been taken to argue in favor of two independent physical processes, which \\citep{Peng:2010fk} called mass-quenching (independent of environment) and environment-quenching (independent of stellar mass),  a terminology which we also adopt in this paper.   These two quenching processes seem to be quite sharp, in the sense that the sSFRs of the surviving star-forming galaxies that are not quenched are almost constant with both galaxy stellar mass and  environment (Peng et al.\\  2010, 2012).  The physical nature of mass- and environment-quenching is however unclear. Several theoretical suggestions have been advanced for what we here are calling mass-quenching.   Some invoke some form of halo physics (e.g., \\citealt{2007MNRAS.380..339B}, Woo et al.\\ 2013, Hearin \\& Watson 2013), and others include AGN feedback (see e.g., Haas et al.\\ 2012 and references therein).  Other possibilities include so-called morphological quenching (\\citealt{2008MNRAS.383..119D}, Martig et al.\\ 2009,2013,  Genzel et al.\\ 2013).   Likewise, for environment-quenching, many physical processes have been proposed, including  ram pressure stripping, accretion shocks, and removal of hot gas (Birnboim \\& Dekel 2003, Kawata \\& Mulchaey 2008, \\citealt{2013arXiv1311.5916C}), with  most explanations generally invoking  some form of gas removal from galaxies as they fall into a group potential (e.g. Feldmann et al.\\ 2011, De Lucia et al.\\ 2012 and references therein).  The difficulty  in identifying the mechanisms that are acting in both quenching mechanisms is at least in part due to both real and spurious correlations between different definitions of mass and environment which blur our understanding of `which mass' and `which environment' are  the most relevant for galaxy evolution. Galaxy stellar mass is directly related to dark matter halo mass, at least for central galaxies. A wide range of environmental measures on different scales may be correlated with each other, from  the global potential of the group dark matter halos, quantified by the halo mass, to the position and thus local density within the halos, to the density provided by  the large scale structure  cosmic web (e.g., \\citealt{van-den-Bosch:2008fk,Kimm:2009ij,Hoyle:2011fk,peng_et_al_2012,2013MNRAS.428.3306W}; see also \\citealt{2013ApJ...776...71C}, hereafter Paper I).  One fact that has become evident, however, is that the rank of a galaxy within the group dark matter halo is very important, i.e., whether the galaxy is the first-ranked central galaxy, lying at the minimum of the potential, or a satellite galaxy orbiting within the group potential  \\citep[e.g.][]{Weinmann:2006vn,Weinmann:2009fk}.  As earlier also explored by \\citet{van-den-Bosch:2008fk},  \\cite{peng_et_al_2012} show that the environmental effects in their 2010 analysis of the SDSS are due to changes in the satellite population alone, i.e. that their 'environment-quenching' is actually a 'satellite-quenching' process.  Within this framework, mass-quenching would be envisioned to act on all galaxies, and be related either directly or indirectly to the stellar mass of the galaxy, whereas environment- (i.e. satellite-) quenching would act only on the satellites.  Since the characteristic $M*$ of centrals and satellites is observed to be very similar, satellites must have experienced the same risk of mass-quenching as central galaxies.  It is an open question however whether this risk occurred when the satellites were actually centrals of their own halos, i.e. before they became satellites of a larger halo (see the discussion in \\citealt{peng_et_al_2012}).  The same authors further showed that the fraction of satellite galaxies with red colors correlates much better with the local over-density, as estimated with a 5th-nearest-neighbour parameter, loosely interpreted to be a measure of location within the groups, than with the optical richness of the group, which is a proxy for the dark matter halo mass (see also Paper I).   Woo et al.\\ (2013) also found that  the fraction of quenched galaxies at fixed galaxy stellar mass shows a  dependence on halo-centric distance only for the satellite population. These authors furthermore reported that the quenching of satellite galaxies depends also on galaxy stellar mass, but only at large (and not at small) halo-centric distances, which they interpret to arise from a quenching dependence on sub-halo mass for satellites that are recently accreted by a more massive halo.  In our picture this could be re-phrased as the increasing importance of mass-dependent mass-quenching  (affecting all galaxies) relative to mass-independent satellite-quenching  at low densities.  Important constraints for establishing the physical mechanisms behind mass-quenching and satellite-quenching can  be obtained by studying the morphologies that are associated with different populations of quenched galaxies.  It has become clear in the last few years that, in addition to the well-known connection between a quenched spectrum and an elliptical morphology, galaxies with a predominant disk component can also be quenched systems \\citep[e.g.][]{Bundy:2010vn,Kovac:2010fk,Oesch:2010zr}. This prompts the question as to whether there is a dependence with environment in the morphological mix of the quenched galaxy population, and especially the quenched satellite population. The quantification of any such dependence may help elucidate in which environments quenching is or is not associated with a morphological transformation, and, more generally, which morphologies and/or morphological transformations can be associated with either mass- or environment-quenching (or both).  In this fourth paper in the  \\emph{Zurich Environmental Study} (ZENS) series we use the ZENS catalog of  galactic and environmental properties published in Paper I to directly address this question by analysing the morphologies of the quenched  {\\it satellite} populations that reside in a set of nearby  galaxy groups that span a range of halo masses and surrounding large-scale structure (LSS) densities.  In ZENS  Paper I, we derived  three  measures of environment, i.e.,  the LSS over-density $\\delta_{LSS}$, the mass of the host group halo $M_{GROUP}$, which we will here rename as $M_{Halo}$,  and the projected halo-centric distance $R$ relative to the characteristic radius of the group halo $R_{200}=\\left(\\frac{GM_{Halo}}{[10H(z)]^2}\\right)^{1/3}$, which we denote $R_{vir}$ (with $H(z)=H_0\\sqrt{\\Omega_M(1+z)^3+\\Omega_{\\Lambda}}$ the Hubble constant at the given redshift).   A strength of ZENS is to  eliminate  sample-dependent biases by using exactly the same sample of galaxies to perform such environment-vs-environment comparisons.  These galaxies are 1455 members of 141 groups which were extracted from the Percolation-Inferred Galaxy Group, 2PIGG, catalog, \\citealt{Eke:2004qf}, of the 2dFGRS, \\citealt{Colless:2001ys}).  The present  paper explores specifically, at fixed galaxy stellar mass; $(i)$ how the fraction of quenched satellites depends on group halo mass, on halo-centric distance and on large-scale structure density,  and $(ii)$ how the morphological composition of the quenched satellite populations depend on these different environmental measures.   Indirect estimates for the quenched fraction of the central galaxies of similar stellar mass are also discussed, in order to obtain a benchmark estimate  for the relative contribution of  mass-quenching and satellite-quenching at a given galaxy mass scale.  The paper is organised as follows. In  Section \\ref{survey}  we briefly describe ZENS, and in particular  the  environmental parameters derived in Paper I  as well as the morphological and spectral  parameters for the ZENS galaxies that are  derived respectively in \\citet{2013ApJ...776...72C}, hereafter Paper II,  and \\citet{2013ApJ...777..116C}, hereafter Paper III\\addtocounter{footnote}{-7}\\footnote{The full ZENS catalog of galaxy and environmental measurements is  published  with Paper I.}.  In Appendix \\ref{allrelaxed} we show the impact of including  unrelaxed groups in the present  analysis, and define the ZENS sample that we use in the present study, which includes both relaxed and unrelaxed groups (see Section \\ref{cleanstuff} for the definitions). In Section \\ref{halosize}, we present the dependence on $M_{Halo}$ and $R/R_{vir}$ of the total fraction of quenched satellites  $f_Q$,  and  of the fraction $f_{Q_{ETG}}$ of quenched satellites that have an early-type galaxy (ETG) morphology, relative to the total quenched satellite population.  In Section \\ref{lss},  we  separately report the analysis of the satellite quenched fraction and its morphological composition as a function of LSS density. In  Section  \\ref{discuss} we investigate the global properties of quenched and star-forming satellites, and of their bulges and disks, and discuss our results first in terms of the apparent morphological changes associated with mass- and environment-quenching, and secondly in comparison  with  traditional semi-analytic models (SAMs) of galaxy formation.  Section \\ref{con} summarises our findings and highlights some open questions that emerge from our analysis.  All our ZENS work is based on a $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ and $h=0.7$ cosmological model. ",
        "conclusions": "\\label{con}   Using the ZENS sample of nearby group halos in the mass range $M_{Halo}\\sim10^{12.2-14.9} M_\\odot$, we have   explored how, in the local Universe, the fraction $f_Q$ of $10^{10-11.5}M_\\odot$ satellites  that are quenched, and the morphological mix of these quenched satellites $f_{Q_{ETG}}$, depend on group halo mass, on normalized halo-centric distance and on the surrounding LSS density.  Similarly to other studies, we find that the total quenched fraction $f_Q$ does not depend much on halo mass or on LSS density.  At the galaxy mass scales that we probe, an average of about 50$\\%$ (at low stellar masses with a median $M_{galaxy}\\sim10^{10.35} M_\\odot$) and $\\sim70\\%$ (at high stellar masses with a median $M_{galaxy}\\sim10^{10.85} M_\\odot$) of satellites have been quenched in one way or the other.   The fraction of quenched satellites increases with decreasing distance from the group centres, from $\\lta40\\%$ at $\\sim1.5$ virial radii, up to $\\gta60\\%$ in the group centres for satellites in the low galaxy  mass bin, and  $\\sim50\\%$ and $\\gta80\\%$ for satellites in the high galaxy mass bin.  These results are broadly consistent with previous findings, e.g. van den Bosch et al.\\ (2008), Peng et al.\\ (2012), Wetzel et al.\\ (2012).  The observed trends in $f_Q$ are consistent with a picture of two independent quenching channels, one that  acts on all galaxies and is independent of the environment of a galaxy (i.e., the mass-quenching channel of Peng et al.\\ 2010), and another one that acts only on satellites and depends on halo-centric radius but not on halo mass  (the environment- or satellite-quenching suggested by the same, and other, authors).  The morphological mix in the quenched satellite fraction, quantified by the fraction $f_{Q_{ETG}}$ of quenched satellites that have an early-type morphology, remains on the other hand strikingly constant across all environments, i.e., we detect no significant variations in $f_{Q_{ETG}}$ with halo mass in the range $\\sim10^{12.2\\rightarrow14.9} M_\\odot$, with halo-centric distance out to and beyond the virial radii of the group halos, or with LSS density (at least  within the range probed by ZENS group with masses $<10^{13.7} M_\\odot$). At low and high galaxy stellar masses, respectively  about 50$\\%$ and 70\\% of  quenched satellites have an early-type morphology in halos of all masses and at any halo-centric radius.  The lack of any dependence on the halo-centric radius, despite the clear change in $f_Q$, shows that the morphological mix of quenched galaxies at a given mass is decoupled from the overall quenched fraction.  The lack of any environmental trend with radius of the morphological mix of the quenched satellite population suggests that both mass- and satellite-quenching processes have the same transformational effect on the morphologies of the galaxies and/or that they have no effect at all. We furthermore note here that the idea of two {\\it independent} quenching processes is argued in the 2010 Peng et al.\\  analysis on the basis of the separability of their red fraction with mass and environment. It is however possible that the underlying physical mechanisms at work in satellite-quenching could be the same as for mass-quenching, i.e. that they are two different manifestations of the same underlying physical process. One way of doing this would be if satellite quenching is the result of a process linked to the parent halo mass.  The striking mass-independence of satellite quenching could then arise if the distribution of host halo masses  for satellites (at the halo masses of interest) was largely independent of the stellar mass of the satellite, which would be supported by the independence of the satellite mass function on group halo mass (e.g., Peng et al.\\ 2012).  Although there appears to be no difference in the morphological outcomes of the two quenching processes, the morphologies of the quenched satellites are not surprisingly systematically different from their star-forming counterparts.  The quenched galaxies have higher $B/T$ and smaller half-light radii.  The morphology-density relation of Dressler (1980) therefore appears to reflect changes in the quenched fraction with radius in the halo rather than changes in the morphological mix of quenched galaxies.  The total $I$-band luminosities and light profiles of the bulges of quenched satellites are very similar to those of the star-forming ones.   Any mass growth in the bulges associated with quenching cannot significantly change their luminosities.  Instead, the changes in $B/T$ and in overall half-light radii are associated with differences in the disks.  Surface brightness fading of previously star-forming disks after star-formation ceases increases the morphological (light-defined) $B/T$ and decreases the half-light radii, even if there are no underlying structural changes in the stellar mass distributions.  We show that such fading can quantitatively explain the morphological differences in the $B/T$ between star-forming and quenched satellites, and also will reduce the half-light radii.  The quenched satellites also exhibit systematically smaller disk scale lengths, which further shrinks the half-light radii relative to those given by uniform disk-fading. These smaller disk scale lengths could be the result either of differential fading with radius in the disks, e.g. due to age gradients, or if the disks of galaxies were systematically smaller, at a given mass, at earlier times.  Both of these would be expected if individual disks grow inside-out.  The overall conclusion is that, at least at the relatively low redshifts being probed here, neither mass-quenching nor satellite-quenching are likely to be changing the mass-defined B/T, which is therefore probably set by other processes operating prior to the onset of quenching.    This emerging picture from the low redshift analysis is consistent with a merger origin for the most massive, pressure-supported spheroids in the local Universe, and disk-fading of bulge-dominated structures as  a key process leading to the typical rotation-supported S0+Elliptical population around $M*$. Of course, the situation at much higher redshifts may or may not be the same.  Current SAMs postulate that early-type morphologies are the result of mergers preceding the infall of the satellite into the virialized group potential, and that subsequent  group halo physics, in the form of gradual removal of galaxies' hot gas halos, leaves the galaxies morphologies unchanged. This picture is not without challenges, not only because the models overpredict the total quenched satellite fraction in all environments, but also because the role of mergers and the quantitative merger rates need to be better defined as a function of cosmic epoch. Specifically, merger rates need to be  made consistent with the emerging paradigm of a non-merging, star-forming galaxy main sequence (e.g., Daddi et al 2007) whose star-formation is fed in quasi-steady-state, from early times, by direct accretion of relatively cold gas streams (e.g., Dekel et al.\\ 2009; see also Dave' et al.\\ 2012 and Lilly et al.\\ 2013). Nevertheless,  SAMs which include a  progressive  removal of the hot gas reservoir of galaxies within group halos well reproduce both the lack of dependence on halo mass and   the dependence on halo-centric distance of the quenched satellite fraction that we see in our ZENS data.  In all environments, they over-predict however  the fraction of quenched satellites with a dominant disk component. Overall, these model successes and failures in  matching the observed quenched satellite  population at redshift zero indicate  that  further model improvements should involve not only the  physical recipes for gas removal from galaxies in heavy halos, but also their star formation and mass assembly histories as a function of galaxy mass."
    },
    "1402/1402.1344_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small Accumulation of dust and ice particles into planetesimals is an important step in the planet formation process. Planetesimals are the seeds of both terrestrial planets and the solid cores of gas and ice giants forming by core accretion. Left-over planetesimals in the form of asteroids, trans-Neptunian objects and comets provide a unique record of the physical conditions in the solar nebula. Debris from planetesimal collisions around other stars signposts that the planetesimal formation process, and hence planet formation, is ubiquitous in the Galaxy. The planetesimal formation stage extends from micrometer-sized dust and ice to bodies which can undergo run-away accretion. The latter ranges in size from 1 km to 1000 km, dependent on the planetesimal eccentricity excited by turbulent gas density fluctuations. Particles face many barriers during this growth, arising mainly from inefficient sticking, fragmentation and radial drift. Two promising growth pathways are mass transfer, where small aggregates transfer up to 50\\% of their mass in high-speed collisions with much larger targets, and fluffy growth, where aggregate cross sections and sticking probabilities are enhanced by a low internal density. A wide range of particle sizes, from mm to 10 m, concentrate in the turbulent gas flow. Overdense filaments fragment gravitationally into bound particle clumps, with most mass entering planetesimals of contracted radii from 100 to 500 km, depending on local disc properties. We propose a hybrid model for planetesimal formation where particle growth starts unaided by self-gravity but later proceeds inside gravitationally collapsing pebble clumps to form planetesimals with a wide range of sizes. \\\\~\\\\~\\\\~}% ",
        "introduction": "}  Most stars are born surrounded by a thin protoplanetary disc with a characteristic mass between 0.01\\% and 10\\% of the mass of the central star \\citep{AndrewsWilliams2005}. Planetesimal formation takes place as the embedded dust and ice particles collide and grow to ever larger bodies. Tiny particles collide gently due to Brownian motion, while larger aggregates achieve higher and higher collision speeds as they gradually decouple from the smallest eddies of the turbulent gas flow and thus no longer inherit the incompressibility of the gas \\citep{Voelk+etal1980}. The gas disc is partially pressure-supported in the radial direction, causing particles of sizes from centimeters to meters to drift towards the star \\citep{Whipple1972,Weidenschilling1977a}. Drift speeds depend on the particle size and hence different-sized particles experience high-speed collisions.  The growth from dust and ice grains to planetesimals -- the seeds of both terrestrial planets as well as the cores of gas giants and ice giants -- is an important step towards the assembly of a planetary system. A planetesimal can be defined as a body which is held together by self-gravity rather than material strength, corresponding to minimum sizes of order 100--1,000 meters \\citep{Benz2000}. The planetesimal formation {\\it stage}, on the other hand, must extend to sizes where the escape speed of the largest bodies exceeds the random motion of the planetesimals, to enter the next stage of gravity-driven collisions. The random speed of planetesimals and preplanetesimals (the latter can be roughly defined as bodies larger than 10 meters in size) is excited mainly by the stochastic gravitational pull from density fluctuations in the turbulent gas; under nominal turbulence conditions the largest planetesimals need to reach sizes as large as 1,000 km to start run-away growth \\citep{Ida+etal2008}. In disc regions with less vigorous turbulence, planetesimal sizes between 1 and 10 kilometers may suffice to achieve significant gravitational focusing \\citep{Gressel+etal2012}.  Planetesimals must grow to run-away sizes despite bouncing and disruptive collisions \\citep{BlumWurm2000}. Planetesimals must also grow rapidly -- radial drift time-scales of cm-m-sized particles are as short as a few hundred orbits. Despite these difficulties there is ample evidence from cosmochemistry that large planetesimals formed in the solar nebula within a few million years \\citep{Kleine+etal2004,Baker+etal2005}, early enough to melt and differentiate by decay of short-lived radionuclides.  The time since the last {\\em Protostars \\& Planets} review on planetesimal formation \\citep{Dominik+etal2007} has seen a large expansion in the complexity and realism of planetesimal formation studies. In this review we focus on three areas in which major progress has been obtained, namely (i) the identification of a bouncing barrier at mm sizes and the possibility of growth by mass transfer in high-speed collisions \\citep{Wurm+etal2005,Johansen+etal2008,Guettler+etal2010,Zsom+etal2010,Windmark+etal2012a}, (ii) numerical simulations of collisions between highly porous ice aggregates which can grow past the bouncing and radial drift barriers due to efficient sticking and cross sections that are greatly enhanced by a low internal density \\citep{Wada+etal2008,Wada+etal2009,SeizingerKley2013}, and (iii) hydrodynamical and magnetohydrodynamical simulations which have identified a number of mechanisms for concentrating particles in the turbulent gas flow of protoplanetary discs, followed by a gravitational fragmentation of the overdense filaments to form planetesimals with characteristic radii larger than 100 km \\citep{Johansen+etal2007,Lyra+etal2008b,Lyra+etal2009a,Johansen+etal2009b,BaiStone2010b,Raettig+etal2013}.  The review is laid out as follows. In section \\ref{s:belts} we give an overview of the observed properties of planetesimal belts around the Sun and other stars. In section \\ref{s:formation} we review the physics of planetesimal formation and derive the planetesimal sizes necessary to progress towards planetary systems. Section \\ref{s:concentration} concerns particle concentration and the density environment in which planetesimals form, while the following section \\ref{s:dustgrowth} describes laboratory experiments to determine the outcome of collisions and methods for solving the coagulation equation. In section \\ref{s:coagulation-fragmentation} we discuss the bouncing barrier for silicate dust and the possibility of mass transfer in high-speed collisions. Section \\ref{s:fluffy} discusses computer simulations of highly porous ice aggregates whose low internal density can lead to high growth rates. Finally in section \\ref{s:unified} we integrate the knowledge gained over recent years to propose a unified model for planetesimal formation involving both coagulation, particle concentration and self-gravity. ",
        "conclusions": ""
    },
    "1402/1402.6108_arXiv.txt": {
        "abstract": "We derive the exact frame-dragging rate inside rotating neutron stars. This formula is applied to show that the frame-dragging rate monotonically decreases from the centre to the surface of the neutron star along the pole. In case of frame-dragging rate along the equatorial distance, it decreases initially away from the centre, becomes negligibly small well before  the surface of the neutron star, rises again and finally approaches to a small value at the surface. The appearance of local maximum and minimum in this case is the result of the dependence of frame-dragging frequency on the distance and angle. Moving from the equator to the pole, it is observed that this local maximum and minimum in the frame-dragging rate along the equator disappears after crossing a critical angle. It is also noted that the positions of local maximum and minimum of the frame-dragging rate along the equator depend on the rotation frequency and central energy density of a particular pulsar. ",
        "introduction": "Compact astrophysical objects such as neutron stars and black holes are the laboratories for the study of Einstein's general relativity in strong gravitational fields. The frame-dragging is one such important general relativistic effect as demonstrated by Lense and Thirring \\cite{LT}. A stationary spacetime with angular momentum shows an effect by which the locally inertial frames are dragged along the rotating spacetime. This makes any test gyroscope in such spacetime precess with a certain frequency called the frame-dragging frequency or the Lense-Thirring (LT) precession frequency $(\\O_{LT})$. The Lense-Thirring frequency is proportional to the angular momentum and compactness of the rotating astrophysical compact object. This effect for a test gyroscope had been calculated and was shown to fall with the inverse cube of the distance of the test gyroscope from the source and vanishes at large enough distances where the curvature effects are small. The precession frequency is thus expected to be larger near the surface of a neutron star and in its interior, rather than at large distance from the star.  The precise mass measurement of PSR J0348+0432 confirmed the existence of a massive neutron star ( $> 2 M_{\\odot}$) \\citep{anto}. It is also known that some of them are observed to possess very high angular velocities. Hence the spacetime curvature would be much higher in the surroundings of those massive neutron stars and the frame dragging effect also becomes very significant in the strong gravitational fields of those rotating neutron stars. It should be noted that the inertial frames are dragged not only outside but also inside the rotating neutron stars.  The theoretical prescription to determine the rate of the frame-dragging precession inside the rotating neutron star was first given by Hartle \\cite{hr67}. In this formalism, one can estimate the frame-dragging precession rate inside a slowly rotating ($\\O R<<c$, where $R$ is the radius of the pulsar, $c$ is the speed of the light in vacuum) neutron star. The final expression of frame-dragging precession rate depends solely on $r$, the distance from the centre of the star,  due to the slow rotation approximation, in Hartle's formalism. It was observed that the frame-dragging frequency was higher at the centre of the star than the frame dragging frequency at the surface. The maximum frame dragging frequency at the centre $(r=0)$ would never exceed the frequency of the rotating neutron star. The frame-dragging effect was applied to various astrophysical problems using Hartle's formalism. Hartle studied this effect on the equilibrium structures of rotating neutron stars \\citep{hr68}. The impact of frame dragging on the Kepler frequency was investigated by Glendenning and Weber \\citep{GW}. It was also demonstrated how this effect might influence the moment of inertia of a rotating neutron star \\citep{web}. Furthermore, Morsink and Stella studied the role of frame dragging in explaining the Quasi Periodic Oscillations of accreting neutron stars \\citep{mor}. Morsink and Stella estimated the precession frequency $\\nu_p$ of the disk's orbital plane about the star's axis of symmetry as the difference between the frequency of oscillations of the particle along the longitude and latitude $(2\\pi\\nu_p=d\\phi/dt-d\\theta/dt)$ observed at infinity. This expression contains the total precession frequency of the disk's orbital plane due to the Lense-Thirring (LT) effect as well as the star's oblateness. Their calculation introduced zero angular momentum observer (ZAMO) and the precession frequency $\\nu_p$ was observed at infinity. It was found that the LT frequency was proportional to the ZAMO frequency on the equatorial plane in the slow rotation limit. This is similar to the Hartle's formalism\\citep{hr67} where the angular velocity $(d\\phi/dt)$ acquired by an observer who falls freely from infinity to the point $(r,\\th)$, is taken as {\\it the rate of rotation of the inertial frame at that point relative to the distant stars.}  In this present manuscript we derive the exact Lense-Thirring precession frequency $(\\Omega_{LT})$ which is measured by a Copernican observer of a gyroscope such as the Gravity Probe B satellite in a realistic orbit\\cite{jh}. In this case, $\\Omega_{LT}$ would not only be the function of $\\omega$ but a complicated function of other metric components also even in the slow rotation limit.  We should note that inside the rapidly rotating stars, one should not a priori expect the similar variation of the precession rates along the equatorial and polar plane. Thus the frame-dragging frequency should depend also on the colatitude ($\\th$) of the position of the test gyroscope. This did not arise in the formalism of Hartle due to the slow-rotation limit. We also note that the LT precession must depend on both the radial distance ($r$) and the colatitude ($\\th$) (see Eq.(14.34) of ~\\cite{jh}) in very weak gravitational fields (far away from the surface of the rotating object).  The exact Lense-Thirring precession rate in strongly curved stationary spacetime had been discussed in detail by Chakraborty and Majumdar~\\cite{cm}. Later, Chakraborty and Pradhan~\\cite{cp} applied this formulation in various stationary and axisymmetric spacetimes. Our main motivation of this paper is to compute the exact LT precession rate inside the rotating neutron star. In this article we avoid all types of approximations and assumptions to obtain the exact LT precession rate inside the rotating neutron stars.  The paper is organized as follows. In section 2 we present the basic equations of frame-dragging effect inside the rotating neutron stars. The numerical method, which has been adopted in the whole paper, is discussed in section 3. We discuss our results in section 4. Finally we conclude in section 5 with a summary. ",
        "conclusions": "\\begin{table} \\begin{center} \\begin{tabular}{|l|l|l|l|l|} \\hline EoS&P(ms)&$\\ve_c(10^{15}$ g/cm$^3$)&$M_G/M_{\\odot}$& $R$ (km)\\\\ \\hline\\hline APR& static & 2.78 & 2.190 & 9.93 \\\\ & 0.6291 & 1.50 & 2.397 & 14.53 \\\\ DD2& static & 1.94 & 2.417 & 11.90 \\\\ & 0.7836 & 1.00 & 2.677 & 17.53 \\\\ Chiral& static & 1.99 & 2.050 & 12.14 \\\\ & 0.8778 & 1.11 & 2.353 & 18.17 \\\\ \\hline \\end{tabular} \\caption{Maximum gravitational masses $(M_G/M_{\\odot})$, equatorial radii $(R)$, and their corresponding central energy densities $(\\ve_c)$ for static $(\\O=0)$ and Keplerian limit $(P=P_K=2\\pi/\\O_K)$ with different EoS, where $P_K$ is the Kepler period in millisecond.} \\label{table_0} \\end{center} \\end{table}  We divide our results into two parts: in the first part we show the frame-dragging effect in some pulsars which rotate with fixed values of Kepler frequencies $\\O_K$ and central densities $\\ve_c$. Next we consider pulsars whose masses and rotational periods are known from observations. Rotational frequencies of observed pulsars are generally much lower than their Kepler frequencies $(\\O<\\O_K)$.  \\subsection{Pulsars rotate with their Kepler frequencies $\\O=\\O_K$}  Figure {\\ref{fig_apr} displays the frame-dragging frequency (or Lense-Thirring precession frequency) as a function of radial distance for APR EoS. Panel (a) of the figure represents the results along the equator whereas panel (b) implies those along the pole. For both panels of Fig. {\\ref{fig_apr}}, we consider rotating neutron stars with central energy densities 5.2435 $\\times 10^{14}$, 6.404 $\\times 10^{14}$ and 7.534 $\\times 10^{14}$ $g / cm^3$ and their corresponding Kepler frequencies are 4000 (online-version: red), 5000 (online-version: green) and 6000 (online-version: blue) $s^{-1}$, respectively whereas masses of the rotating compact stars in three cases range from $\\sim$ 0.6 to $\\sim$ 1.7 M$_{\\odot}$. The Kepler periods $P_K=$1.57 ms, 1.26 ms and 1.05 ms correspond to the above Kepler frequencies. Here and throughout the paper, $\\O_{LT}$ is measured in a Copernican frame. For the cases in panel (a), frame-dragging frequencies decrease initially with increasing distance from the centre and encounters a local minimum at a distance $r_{min} \\sim 0.5 r_e$ which is well below the surface. It is interesting to note here that the frame-dragging frequencies in all three cases rise again and attain a local maximum at the distance $r_{max} \\sim 0.7 r_e$ and finally drop to smaller values at the surface. On the other hand, the frame-dragging frequencies along the pole smoothly vary from large values at the centre to smaller values at the surface as evident from panel (b) of Fig. {\\ref{fig_apr}}.  Figure {\\ref{fig_dd2}} shows the frame-dragging frequency along the equator (panel (a)) and along the pole (panel (b)) for the DD2 EoS whereas Figure {\\ref{fig_chi}} exhibits the frame-dragging frequency along the the equatorial distance (panel (a)) and polar distance (panel (b)) for the Chiral EoS. In both figures, results are shown for Keplerian frequencies $\\O_K =$ 4000, 5000 and 6000 $s^{-1}$. However, the central energy densities corresponding to the Keplerian frequencies mentioned above are different for three EoSs. The behaviour of frame-dragging frequencies along the equator and pole in Fig. {\\ref{fig_dd2}} and Fig. {\\ref{fig_chi}} is qualitatively similar to the results of Fig. {\\ref{fig_apr}}. In Figs. {\\ref{fig_apr}} - {\\ref{fig_chi}}, as the rotation frequency $\\O_K$ and the central energy density $\\ve_c$ increase, the frame-dragging frequencies increase and also the local maxima and minima shift towards the surface of the neutron star along the equator for all three EoSs. It reveals an important conclusion that {\\it the ratio of the positions of the local maxima and minima to the radius of the neutron star must depend on $\\O$ and $\\ve_c$ for a particular pulsar.}  It could be easily seen from Eq. (\\ref{ltn}) that the Lense-Thirring frequency inside a neutron star is a function of both the radial distance $r$ and colatitude $\\th$. The colatitude plays a major role to determine the exact frame-dragging frequency at a particular point inside the rotating neutron star as evident from Figs. {\\ref{fig_apr} - {\\ref{fig_chi}. We obtain the Lense-Thirring frequency at the pole by just plugging-in $\\th = 0$ in Eq. (\\ref{ltn}) and it is given by $\\O_{LT} = e^{-(\\alpha+\\sigma)}\\omega$. It should be noted here that the Lense-Thirring frequency is connected to $\\omega$ which appears as the non vanishing metric component in the metric of the rotating star. According to the theorem by Hartle, the dragging of inertial frames as represented by $\\omega$ with respect to a distant observer decreases smoothly as a function of $r$ from a large value at the centre of the star to a smaller value at the surface \\citep{web} for both equatorial and polar cases. In this formalism the frame-dragging frequency depends solely on $r$. For a fixed value of $r$, one gets the same frequency from the equator to the pole inside the rotating neutron star. We obtain the similar behaviour of the Lense-Thirring frequency along the pole in panel (b) of Figs. {\\ref{fig_apr}} - {\\ref{fig_chi}} as obtained in Hartle's formalism.} However, our results along the equator are quite different from what was obtained using Hartle's formalism \\citep{web}. It is evident from Figs. \\ref{fig_apr}-\\ref{fig_chi} that the plots are smooth along the pole but not along the equator. For the calculation of the Lense-Thirring frequency along the equator, we find that the second term of Eq. (\\ref{ltn}) does not contribute. Further investigation of the first term involving metric components $\\sigma$, $\\omega$ and their derivatives reveals that this term is responsible for the local maxima and minima along the equator as reported above. The appearance of local maxima and minima in the Lense-Thirring frequencies along the equator may be attributed to the dependence of $\\O_{LT}$ on $r$ and $\\th$. As a consistency check, we obtain two solutions for local maximum and minimum after extremising the Eq. ({\\ref{ltn}}) with respect to $r$. Details are given by {Appendix \\ref{apndx_2}}.  \\begin{figure} \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{apr_eq3.eps}} \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{apr_pole3.eps}} \\caption{\\label{fig_apr}\\textit{Frame-dragging effect inside the rotating neutron stars from the origin to the surface, calculated for the APR EoS. $\\O_K$ and $\\ve_c$ denote the Kepler frequency and the central star density, respectively. Surface of the neutron star along the pole located around $0.6r_e$ but the plot is still valid beyond the surface of the pole as our formalism is applicable for regions outside the pulsar.}} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{dd2_eq3.eps}} \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{dd2_pole3.eps}} \\caption{\\label {fig_dd2}{\\textit{Same as Fig. \\ref{fig_apr}, but calculated for the DD2 EoS}}} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{CHIRAL_eq3.eps}} \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{CHIRAL_pole3.eps}} \\caption{\\label{fig_chi}{\\textit{Same as Fig. \\ref{fig_apr}, but calculated for the Chiral EoS.}}} \\end{center} \\end{figure}  The normalised frame-dragging values at the centre $(\\ti{\\O}_c = {\\O_{LT}^{centre}}/{\\O})$ and the surface $(\\ti{\\O}_s = {\\O_{LT}^{surface}}/{\\O})$ of the star models with three EoSs are recorded in Table {\\ref{table_1}}. \\begin{table} \\begin{center} \\begin{tabular}{|l|l||l|l|l||l|l|l|} \\hline &$P_K$&\\multicolumn{3}{l|}{Along the equator} &\\multicolumn{3}{l|}{Along the pole}\\\\ \\cline{3-8} &(ms)&APR&DD2&Chiral&APR&DD2&Chiral\\\\ \\hline\\hline & 1.57 & 0.008 & 0.013 & 0.019 & 0.046 & 0.069 & 0.099 \\\\ $\\ti{\\O}_s$& 1.26 & 0.016 & 0.029 & 0.040 & 0.087 & 0.139 & 0.184 \\\\ & 1.05 & 0.031 & 0.059 & 0.070 & 0.151 & 0.252 & 0.287 \\\\ \\hline & 1.57 & 0.242 & 0.286 & 0.356 & 0.242 & 0.286 & 0.356 \\\\ $\\ti{\\O}_c$ & 1.26 & 0.354 & 0.457 & 0.580 & 0.354 & 0.457 & 0.580 \\\\ & 1.05 & 0.515 & 0.723 & 0.884 & 0.515 & 0.723 & 0.884 \\\\ \\hline \\end{tabular} \\caption{Normalised angular velocities of the local inertial frame-dragging at the surface $\\ti{\\O}_s$ and the centre $\\ti{\\O}_c$  of the neutron stars which are rotating at their respective Kepler periods $(P_K\\equiv 2\\pi/\\O_K)$ as measured by a distant observer.} \\label{table_1} \\end{center} \\end{table} It is noted that the normalised frame-dragging value at the star's center is maximum and falls off on the surfaces of the equator and pole for three EoSs irrespective of whether the compact star is rotating slowly or fast. However, for a particular EoS, the normalised frame-dragging value at the star's centre and surface is higher for a fast rotating star with $P_K = 1.05$ ms than those of a slowly rotating star with $P_K = 1.57$ ms for both cases along the equator and pole. One can see another interesting thing from the Table \\ref{table_1} that $\\ti{\\O}_s$ is always higher at the pole than $\\ti{\\O}_s$ at the equator for a particular pulsar. It is due to the effect of rotation frequency $\\O$ (of the star) for which pole is nearer to the center than the surface as it is evident from Table \\ref{table_1}. Thus, the inertial frame-dragging effect is higher at the surface of the neutron star along the pole than that at the surface of the neutron star along the equator.  Now we investigate the dependence of local maxima and minima in $\\O_{LT}$ along the equator on the angle $\\th$. In Figure {\\ref{cr}}, $\\O_{LT}$ is shown as a function of $s$ defined by Eq. (\\ref{rs}) and $\\cos\\th$ for DD2 (panel(a)) and Chiral (panel (b)) EoSs. The local maxima and minima in $\\O_{LT}$ along the equator are clearly visible in both panels of Fig. {\\ref{cr}}. It has been already noted that $\\O_{LT}$ along the pole decreases smoothly from the center to the surface. As the latitude ($\\th'=\\pi/2-\\th$) increases from the equator to the pole, the height between the maximum and minimum of $\\O_{LT}$ diminishes and after a certain `critical' angle ($\\th_{cr}$) both extrema disappear and the plot is smooth like the plot along the pole. The critical angle could be seen from the 3-D plot in Fig. \\ref{cr}. \\begin{figure} \\begin{center} \\subfigure[$DD$2 EoS]{ \\includegraphics[width=2in,angle=-90]{dd2_kep5000.eps}} \\subfigure[APR EoS]{ \\includegraphics[width=2in,angle=-90]{apr_kep5000.eps}} \\caption{\\label{cr}{\\textit{3-D plots of $\\O_{LT}$ of the pulsar which is rotating with $\\O_K=5000\\, \\text{s}^{-1}$} as a function of $s$ and $\\cos\\th$ for (a) DD2 EoS and (b) APR EoS.}} \\end{center} \\end{figure} This value of the critical angle is $\\mu\\approx 0.5$ or $\\th_{cr}'=30^{\\circ}$ where local maximum  $(r_{max})$ and minimum $(r_{min})$ disappear. So, if we plot $\\O_{LT}$ vs $\\th$ at the point $r_{max}$ for a specific Kepler frequency, (namely $\\O_K=5000\\, \\text{s}^{-1}$), we could find that the frame-dragging frequency increases from the equator to the pole for the specific $r_{max}$ as it is exhibited by Fig. {\\ref{cr}}.  \\subsection{Pulsars rotate with their frequencies $\\O<\\O_K$}  Now we apply our exact formula of $\\O_{LT}$ to three known pulsars. Three pulsars chosen for this purpose are J1807-2500B, J0737-3039A and B1257+12. Periods of those pulsars are given by Table {\\ref{table_2}}. Masses of those pulsars are also known and range from 1.337 to 1.5 M$_{\\odot}$. Furthermore, we adopt the same EoSs in this calculation as considered in the previous sub-section. Though periods of these pulsars are larger than the Keplerian periods, the calculation of $\\O_{LT}$ inside these real pulsars are equally important like the cases with Kepler frequencies demonstrated already.  We calculate the normalised angular velocity at the centre and surface of these pulsars with APR, DD2 and Chiral EoSs. It is noted from Table \\ref{table_2} that the behaviour of the normalised angular velocity from the centre to the surface or along the pole and equator is qualitatively same as shown in Table {\\ref{table_1}}. \\begin{table} \\begin{center} \\begin{tabular}{|l|l|l||l|l|l||l|l|l|} \\hline &Name of&$P$&\\multicolumn{3}{l|}{Along the equator} &\\multicolumn{3}{l|}{Along the pole}\\\\ \\cline{4-9} &the Pulsar&(ms)&APR&DD2&Chiral&APR&DD2&Chiral\\\\ \\hline\\hline & J1807-2500B & 4.19 & 0.099 & 0.075 & 0.064 & 0.156 & 0.120 & 0.105 \\\\ $\\ti{\\O}_s$& J0737-3039A & 22.70 & 0.095 & 0.073 & 0.062 & 0.154 & 0.122 & 0.106 \\\\ & B1257+12 & 6.22 & 0.122 & 0.091 & 0.077 & 0.188 & 0.145 & 0.126 \\\\ \\hline & J1807-2500B & 4.19 & 0.707 & 0.548 & 0.516 & 0.707 & 0.548 & 0.516 \\\\ $\\ti{\\O}_c$ & J0737-3039A & 22.70& 0.685 & 0.538 & 0.502 & 0.685 & 0.538 & 0.502 \\\\ & B1257+12 & 6.22 & 0.825 & 0.632 & 0.601 & 0.825 & 0.632 & 0.601 \\\\ \\hline \\end{tabular} \\caption{Normalised angular velocities of the local inertial frame-dragging at the surface $\\ti{\\O}_s$ and the centre $\\ti{\\O}_c$ of some known rotating neutron stars.} \\label{table_2} \\end{center} \\end{table}  We also plot the frame-dragging frequency ($\\O_{LT}$) as a function of radial distance along the equator (panel (a)) and pole (panel (b)) for three EoSs in Figures {\\ref{j1807}} - {\\ref{b1257}}. The frame-dragging frequency behaves smoothly along the pole from the centre to the surface as shown by panel (b) of these figures. Results of panel (a) of the figures show similar features of local maxima and minima along the equator as found in Figs. \\ref{fig_apr}-\\ref{fig_chi}. We note that all the local minima of $\\O_{LT}$ are located around $r_{min}\\sim 0.7r_e$ and the local maxima are located around $r_{max}\\sim 0.9r_e$ in Figs. {\\ref{j1807}} - {\\ref{b1257}}. \\begin{figure}[h!] \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{j1807_eq.eps} } \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{j1807_pole.eps} } \\caption{\\label{j1807} \\textit{Frame-dragging effect inside the rotating neutron star from the origin to the surface, calculated for J1807-2500B ($M=1.366 M_{\\odot}, \\O = 1500.935\\, \\text{s}^{-1}$)}} \\end{center} \\end{figure}  \\begin{figure}[h!] \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{j0737_eq.eps} } \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{j0737_pole.eps} } \\caption{\\label{j0737} \\textit{Frame-dragging effect inside the rotating neutron star from the origin to the surface, calculated for J0737-3039A ($M=1.337 M_{\\odot}, \\O = 276.8\\, \\text{s}^{-1}$)}} \\end{center} \\end{figure}  \\begin{figure}[h!] \\begin{center} \\subfigure[along the equator]{ \\includegraphics[width=2in,angle=-90]{b1257_eq.eps} } \\subfigure[along the pole]{ \\includegraphics[width=2in,angle=-90]{b1257_pole.eps} } \\caption{\\label{b1257} \\textit{Frame-dragging effect inside the rotating neutron star from the origin to the surface, calculated for B1257+12 ($M=1.5 M_{\\odot}, \\O = 1010.321\\, \\text{s}^{-1}$)}} \\end{center} \\end{figure} We also plot the $\\O_{LT}$ of pulsar J0737-3039A as a function of $s$ and $\\cos\\th$ for DD2 (panel (a)) and APR (panel (b)) EoSs in Figure {\\ref{cr2}. It is noted from Fig. \\ref{cr2} that the value of $\\th_{cr}$ is around $30^{\\circ}$ for the pulsar J0737-3039A for DD2 and APR EoSs.  \\begin{figure} \\begin{center} \\subfigure[$DD$2 EoS]{ \\includegraphics[width=2in,angle=-90]{dd2_j0737.eps}} \\subfigure[APR EoS]{ \\includegraphics[width=2in,angle=-90]{apr_j0737.eps}} \\caption{\\label{cr2}\\textit{3-D plots of $\\O_{LT}$ of the pulsar J0737-3039A as a function of $s$ and $\\cos\\th$ for DD2 (panel (a)) and APR (panel (b)) EoSs.}} \\end{center} \\end{figure}"
    },
    "1402/1402.1034_arXiv.txt": {
        "abstract": "{} {Following the recent discovery of a large population of young stars in front of the Orion Nebula, we carried out an observational campaign with the \\emph{DECam} wide-field camera covering $\\approx$10~deg$^{2}$ centered on NGC 1980 to confirm, probe the extent of, and characterize this foreground population of pre-main-sequence stars.} {We used  multiwavelength wide-field images and catalogs to identify potential foreground pre-main-sequence stars using a novel probabilistic technique based on a careful selection of colors and luminosities.} {We confirm the presence of a large foreground population towards the Orion A cloud. This population contains several distinct subgroups, including NGC~1980 and NGC~1981, and stretches across several degrees in front of the Orion~A cloud. By comparing the location of their sequence in various color-magnitude diagrams with other clusters, we found a distance and an age of 380~pc and 5$\\sim$10~Myr, in good agreement with previous estimates. Our final sample includes 2123 candidate members and is complete from below the hydrogen-burning limit to about 0.3~M$_{\\sun}$ , where the data start to be limited by saturation. Extrapolating the mass function to the high masses, we estimate a total number of $\\approx$2\\,600 members in the surveyed region.  } {We confirm the presence of a rich, contiguous, and essentially coeval population of about 2600 foreground stars in front of the Orion~A cloud, loosely clustered around NGC~1980,  NGC~1981, and a new group in the foreground of the OMC-2/3. For the area of the cloud surveyed, this result implies that there are more young stars in the foreground population than young stars inside the cloud. Assuming a normal initial mass function,  we estimate that between one to a few supernovae must have exploded in the foreground population in the past few million years, close to the surface of Orion~A, which might be responsible, together with stellar winds, for the structure and star formation activity in these clouds. This long-overlooked foreground stellar population is of great significance, calling for a revision of the star formation history in this region of the Galaxy.} ",
        "introduction": "\\label{sec:introduction} The Orion star-forming complex is the nearest active star-forming region to Earth that produces massive stars, and is one of the richest in the nearest 1-kpc. It has long been recognized as a benchmark laboratory for star and planet formation studies as well as for the formation and dispersal of OB associations. The entire Orion star formation complex spreads across 200 pc and has spawned about 10$^4$ stars in the past 12 Myr \\citep[e.g.][]{1994A&A...289..101B,2008hsf1.book..459B,2008hsf1.book..483M,2008hsf1.book..838B}. There are at least four subpopulations of young stars within the complex, first identified by \\citet{1964ARA&A...2..213B}, from the older and dust-free Orion OB 1a group in the North (age $\\sim$10 Myr) to the still-embedded Orion OB 1d (age $\\sim$1 Myr), or the Orion Nebula Cluster, to the South. It was realized early that the groups Orion OB 1c and 1d overlapped, at least in part, along the line of sight \\citep[e.g.][]{1978ApJS...36..497W,1998AJ....115.1524G}.  Recently, \\citet{2012A&A...547A..97A} used the denser regions of the Orion A cloud to block the background light in the optical, effectively isolating the stellar population in front of it, and found a rich stellar population in front of this cloud. The surprising result in this paper was not the confirmation of the existence of a foreground population, but how rich this population is, in particular towards the previously unstudied NGC 1980 cluster, which was shown to be slightly older ($\\sim$ 5 Myr) than the embedded population in the cloud and contains a full stellar spectrum from OB stars to substellar objects. Using X-ray observations, \\citet{2013ApJ...768...99P} recently found a population of sources towards the Orion A cloud that shows a very low amount of extinction (on the order of N$_{H} = 3\\times10^{20}$ cm$^{-2}$, or about 0.1~mag of visual extinction) towards $\\iota-$Ori and NGC~1980. This very low extinction population coincides with the foreground population isolated in \\citet{2012A&A...547A..97A}, further confirming its presence.  The existence of a rich foreground population implies that the closest massive star formation and cluster formation benchmark, namely the Orion Nebula Cluster \\citep[e.g.][]{2012ApJ...748...14D}, suffers from significant foreground contamination (on the order of 10\\% for the Orion Nebula Cluster, but higher for the less rich L1641 young population). This result also implies that massive star-forming regions  can have complex star formation histories, and superposition of distinct young populations along a given line of sight probably is the norm, which complicates the determinations of, for example, the stellar yield of a molecular cloud, the duration of the star formation process (age spreads), and mass functions.  \\citet{2012A&A...547A..97A} made a first estimation of the size of the population (about 2000 objects) by extrapolating from the relatively high extinction regions of Orion A to a larger field.  This estimate was necessarily approximate, and called for  a better characterization of the foreground population. In the current paper we make use of data from a dedicated observational campaign with the DECam wide-field camera, covering about 10 square degrees of sky centered on NGC~1980, to confirm, probe the extent of, and characterize the foreground population.  This paper is structured as follows: in Sect. 2 we describe the observational data acquired for this project as well as the archival data used. In Sect. 3 we present the results of our approach, namely the identification of the two foreground populations and their characterization. We present a general discussion on the importance of the result found in Sect. 4 and summarize our main results in Sect. 5. ",
        "conclusions": "We combined new wide-field optical observations of the Orion~A group of associations with archival data and catalogs. A dense sequence appears clearly in all color-magnitude diagrams. The observed match with empirical sequences suggests an age of 5$\\sim$10~Myr and a distance of $\\approx$380~pc, placing the corresponding sources in the immediate front of the Orion A cloud and the ONC. Members of the sequence were selected using a novel multi-dimensional probabilistic analysis using carefully selected color-magnitude diagrams and luminosities. The sources cluster into several groups, which can be spatially associated to NGC~1980, L1641N, L1641W, and NGC~1981, and a new group OriA-Fore~1 located southwest of OMC3. The current work confirms the results reported in \\citet{2012A&A...547A..97A} using X-ray, optical, and infrared data.  We propose the following scenario to explain the observations in the region: a first generation of clusters and associations located at 380$\\sim$400~pc formed between 5$\\sim$10~Myr ago. It included NGC~1980, NGC~1981, 25-Ori, $\\sigma-$Ori, $\\lambda-$Ori, and possibly L1641W and OriA-Fore~1.  These groups formed several thousands of stars, including a few massive stars. After a few million years, the most massive stars must have exploded as supernovae, which together with stellar winds very likely triggered episodic star formation events in the ONC, resulting in the Trapezium, L1641N, OMC2/3, and maybe even NGC~2024 in Orion B."
    },
    "1402/1402.4427_arXiv.txt": {
        "abstract": "Between 1997 and 2004 several observing runs were conducted mainly with the CTIO 0.9 m to image ICRF counterparts (mostly QSOs) in order to determine accurate optical positions. Contemporary to these deep CCD images the same fields were observed with the US Naval Observatory (USNO) astrograph in the same bandpass. They provide accurate positions on the Hipparcos/Tycho-2 system for stars in the 10 to 16 magnitude range used as reference stars for the deep CCD imaging data. Here we present final optical position results of 413 sources based on reference stars obtained by dedicated astrograph observations which were reduced following 2 different procedures. These optical positions are compared to radio VLBI positions. The current optical system is not perfectly aligned to the ICRF radio system with rigid body rotation angles of 3 to 5 mas (= 3 $\\sigma$ level) found between them for all 3 axes. Furthermore, statistically, the optical$-$radio position differences are found to exceed the total, combined, known errors in the observations. Systematic errors in the optical reference star positions as well as physical offsets between the centers of optical and radio emissions are both identified as likely causes.  A detrimental, astrophysical, random noise (DARN) component is postulated to be on about the 10 mas level. If confirmed by future observations, this could severely limit the Gaia to ICRF reference frame alignment accuracy to an error of about 0.5 mas per coordinate axis with the current number of sources envisioned to provide the link. A list of 36 ICRF sources without the detection of an optical counterpart to a limiting magnitude of about R=22 is provided as well. ",
        "introduction": "The Hipparcos Catalogue \\citep{hipcat} is currently the primary optical reference frame. It is linked to the defining radio International Celestial Reference Frame (ICRF) by various methods with a dozen radio stars observed with Very Long Baseline Interferometry (VLBI) having the greatest weight \\citep{k97}. An improved version of the ICRF, the ICRF2 \\citep{icrf2} now contains over 3400 compact, extragalactic sources on the same coordinate system, which is assumed to be inertial.  The Hipparcos Celestial Reference frame (HCRF) is entirely based on space mission observations and contains positions, proper motions and parallaxes of over 100,000 bright stars on the milli-arcsecond (mas) level (mas/yr for individual proper motions).  Zonal and magnitude dependent errors in the Hipparcos Catalogue are estimated to be even smaller. The HCRF is estimated to be accurate to 0.6 mas (1$\\sigma$) in its alignment to the ICRF at the epoch of 1991.25 with an estimated 0.25 mas/yr (1$\\sigma$) residual system rotation error per rigid body angle as compared to the ICRF \\citep{k97}.  The Tycho-2 catalog is the first step of the densification of the optical reference frame to contain the 2.5 million brightest stars to visual magnitude about 11 \\citep{tycho2}.  Random position errors of Tycho-2 stars are about 7 mas per coordinate at V = 9 and epoch 1991, increasing for the fainter stars.  The random error of Tycho-2 star proper motions is typically about 2 mas/yr.  The Tycho-2 data are based on Hipparcos star tracker observations (at mean epoch of 1991.25) combined with over 140 ground-based star catalogs to provide sufficient leverage in epoch difference for precise proper motions.  The largest weight of the early epoch data in Tycho-2 comes from the Astrographic Catalogue (AC) \\citep{ac2000}, with positions of photographic plates at an epoch around 1900.  The AC plates cover an area of just over 2$^{\\circ}$ by 2$^{\\circ}$ and Hipparcos stars were used for their astrometric reductions. Although great care has been taken in these reductions, remaining systematic errors on the AC plates are estimated to be on the 100 mas level, maybe more (Urban, private com.).  This results in possible systematic errors of Tycho-2 proper motions on the 1 to 2 mas/yr level.  The goal of this paper is to examine the link between the defining radio reference frame (ICRF) and the optical reference frame (HCRF) by observing the optical counterparts of ICRF sources on the HCRF using the Tycho-2 as reference star catalog in a 2-step process to bridge the large magnitude difference.  While the deep CCD imaging captures the optical counterparts, the USNO astrograph is used to provide intermediate bright (12 to 16 mag) reference stars.  Besides possible systematic errors in the optical reference frame (i.e.~Tycho-2) the assumption of coinciding centroids of the radio and optical centers of emissions of the extragalactic sources is questioned.  Observed optical$-$radio position offsets are found to be larger than expected from the combination of all known, random errors.  Clues for both cases, reference frame issues, and physical offset between radio and optical centers of emission are found.  If the later turns out to be a wide-spread issue on the several mas level this would have serious consequences for the alignment accuracy of the future optical, celestial reference frame from the Gaia mission \\citep{gaia} to the ICRF. Our data are not precise enough to make that case; however, some individual sources are found with relatively large radio-optical position offsets likely on the 30 mas level, a few with even much larger offsets, indicating a physical non-coincidence of the radio-optical centroids.  A possible physical offset between the radio and optical centroid of ICRF sources was first discussed in connection with a correlation to the radio structure index \\citep{dasilva}.  Results from a similar, extensive investigation of 300 optical counterparts were recently published by the Rio de Janeiro group \\citep{riosurvey}. Direct differences between ICRF2 radio positions and optical counterpart positions found in the Sloan Digital Sky Survey (SDSS) DR9 were studied for 1297 common sources in another recent paper \\citep{orosz}. Positions of 22 ICRF sources observed with the SOAR 4.2m and ESO 2.2m telescopes again show a correlation of optical$-$radio position offsets with radio structure index \\citep{camargo}.  Our research presented here is the largest and most accurate such sample published so far related to the original ICRF sources and can be considered as a follow-up to an earlier work based on photographic plates for secondary reference stars \\citep{rorf99}. Preliminary results of a subset of these data in collaboration with another group was published earlier \\citep{a2003}. The high precision of our current investigation could be achieved by using dedicated astrometric observations along with contemporaneous astrograph observations, thus eliminating the degrading effect of proper motion errors of intermediate bright link stars. Partial, preliminary results of our research were presented at several meetings, \\citep{adela}, \\citep{iau2003}, \\citep{lowel2005}, \\citep{iaus248}, \\citep{iau2009} and \\citep{iau2012}. It was not until recently that the final (4th) USNO CCD Astrograph Catalog (UCAC4) was published \\citep{ucac4} and its reduction process could be adopted for the critical astrogaph data to finally conclude with this paper a research program that lasted more than a decade. ",
        "conclusions": "  \\begin{enumerate} \\item This investigation is the most precise of its kind so far with a formal mean error of 25 mas per coordinate for the optical positions of ICRF counterparts. This was achieved from well calibrated astrometric observations with mostly the 0.9 m CTIO telescope in combination with contemporary astrograph observations providing intermediate bright reference stars on the Tycho-2 (Hipparcos) coordinate system (HCRS) at the same epoch as the deep imaging data.  \\item This sets an upper limit of about 25 mas (1$\\sigma$) of possible non-coincidence of the optical and radio centers of emissions of typical ICRF quasars, with a few outliers already seen well exceeding this level of positional precision.  \\item Comparing UCAC2-type and UCAC4-type solutions, use of UCAC2-type reference stars result in slightly smaller scatter of the optical$-$radio position differences, while use of UCAC4-type reference stars show smaller offsets to the radio reference frame coordinate axes.  \\item Larger than expected (from the combined, total, formal errors) optical$-$radio position offsets are seen, which in part can be explained by systematic errors in the reference stars, particularly errors in the Tycho-2 catalog changing on degree scales, likely due to limitations in calibrating magnitude equations and coma terms in the Astrographic Catalogue data.  \\item Correlations of optical$-$radio position differences with redshift and radio structure index of sources, comparison of the number of ``outliers\" from optical surveys of different precisions, and examining Hipparcos to ICRF link observations by HST FGS involving QSO-star pairs indicate other contributions besides systematic errors in reference star positions. These indications lead to the hypothesis of a true, detrimental, astrophysical random noise (DARN) offset of radio and optical centers of emission on about the 10 mas level. Higher optical accuracy data is needed to conclusively prove this DARN hyphothesis; however, besides the observational indications there is a possible astrophysical explanation for such an error contribution on that level due to host galaxy structure at optical wavelengths.  \\item It is estimated that Gaia observations will provide a rigid, instrumental, optical reference frame on the few micro-arcsecond level \\citep{gaiasw}. Its construction will not be affected by possible significant optical$-$radio position offsets of individual ICRF sources. However, the errors in alignment of the Gaia internal coordinate system to the ICRF will be affected by DARN, if it exists.  Although the radio positions as well as the Gaia observed optical positions of ICRF sources in common will have sub-mas precision, DARN will dominate the error budget. If the mean optical$-$radio position offsets are indeed typically 10 mas and there are about 600 link sources (2 coordinates each) the 3 orientation angles of the Gaia frame each can only be determined to about 0.5 mas (1$\\sigma$) precision with respect to the ICRF. Absolute proper motions and parallaxes will not be compromised, only absolute positions, if a rigid, instrumental Gaia system can be constructed without zonal errors.  \\item  If DARN exists and a higher positional alignment of the Gaia and ICRF reference frames is desired, more objects in common with sufficient radio and optical flux will be needed, instead of concentrating on a ``clean\" sample (i.e.~small radio structure index, low optical variability), because the error contribution of an individual source for the radio to optical coordinate system link will be dominated by that unknown, random optical position offset. Ground-based observations of ICRF sources too faint to be seen with Gaia directly could provide optical positions on the Gaia system using Gaia reference stars to its limiting magnitude in small fields of view, using large aperture, long-focus telescopes. Determining radio positions of more targets observed by Gaia is another option.  Radio or optical positions of those additional source observations would need to be only on or slightly better than the position error level of DARN (maybe 5 to 10 mas), to provide meaningful, additional data to improve the link of Gaia and ICRF reference frames. Monitoring programs to arrive at a ``clean\" sample would be of lower priority than the shear number of ``reasonably good\" sources with both optical and radio positional data. The stability of an optical or radio position (below the DARN level) would not matter because of the unknown significant position offset between the centers of emission of optical and radio source. \\end{enumerate}"
    },
    "1402/1402.1202_arXiv.txt": {
        "abstract": " ",
        "introduction": "CCAT \\citep{2012SPIE.8444E..2MW,2013AAS...22115006G,P10_adassxxiii} is a 25\\,m diameter sub-millimeter telescope to be built on Cerro Chajnantor in Chile near the ALMA site at an altitude of 5600\\,m. The additional height above the ALMA array results in significant improvements in transparency across all observing bands (Fig.\\ \\ref{fig:trans}), and 1.64 times better than the ALMA site at 350\\,\\micron\\ \\citep{2011RMxAC..41...87R}. CCAT will initially operate from 350\\micron\\ to 2\\,mm but will be capable in the future of operating at 200\\micron\\ in the very best weather.  The CCAT project has identified four major first generation instruments to achieve its science goals. SWCam \\citep{2013AAS...22115007S} will be the first-light camera having of order 60,000 detectors operating mainly at 350\\micron\\ with additional detectors out to 2\\,mm. CHAI \\citep{GoldsmithCHAI2012} will be a large format heterodyne array operating in two bands with the backend able to process spectra with a bandwidth of 4\\,GHz and 64,000 channels. LWCam \\citep{2013AAS...22115008G} is a dedicated long-wave camera operating in 5-6 bands between 750\\,\\micron\\ and 2.1\\,mm with a long-wavelength goal of 3.3\\,mm. X-Spec \\citep{2013AAS...22115009B} is a multi-object spectrometer with $\\sim$\\,100 beams on the sky, each covering a frequency range of 190-520\\,GHz in two bands simultaneously with a resolving power of 400 -- 700.  \\begin{figure} \\includegraphics[width=\\textwidth]{jenness-hotwired3_f1} \\caption{Atmospheric transmission for exceptional (10\\%), excellent (25\\%), and median conditions at CCAT and for median conditions at ALMA and at Mauna Kea \\citep[ATM model;][]{2001ITAP...49.1683P}. Water vapor (PWV) distributions determined from 350\\micron\\ tipper measurements \\citep{2011RMxAC..41...87R}. (figure credit: S.\\ Radford).} \\label{fig:trans} \\end{figure} ",
        "conclusions": ""
    },
    "1402/1402.6614_arXiv.txt": {
        "abstract": "Redshift drift provides a direct kinematic measurement of cosmic acceleration but it occurs with a characteristic time scale of a Hubble time.  Thus redshift observations with a challenging precision of $10^{-9}$ require a 10 year time span to obtain a signal-to-noise of 1. We discuss theoretical and experimental approaches to address this challenge, potentially requiring less observer time and having greater immunity to common systematics.  On the theoretical side we explore allowing the universe, rather than the observer, to provide long time spans; speculative methods include radial baryon acoustic oscillations, cosmic pulsars, and strongly lensed quasars. On the experimental side, we explore beating down the redshift precision using differential interferometric techniques, including externally dispersed interferometers and spatial heterodyne spectroscopy. Low-redshift emission line galaxies are identified as having high cosmology leverage and systematics control, with an 8 hour exposure on a 10-meter telescope (1000 hours of exposure on a 40-meter telescope) potentially capable of measuring the redshift of a galaxy to a precision of $10^{-8}$ (few $\\times 10^{-10}$). Low-redshift redshift drift also has very strong complementarity with cosmic microwave background measurements, with the combination achieving a dark energy figure of merit of nearly 300 (1400) for 5\\% (1\\%) precision on drift. ",
        "introduction": "\\label{Sec:intro}  Our universe is dynamic, i.e.\\ the metric evolves, or more simply the scale factor of the universe changes with time: hence $a(t)$.  This gives rise to the cosmic redshift of light from distant sources.  Moreover, since the expansion rate itself evolves, in all but a coasting, Milne universe, the redshift of an object with fixed comoving coordinate will shift.  This redshift drift was introduced by \\citet{sandage} and McVittie \\citet{mcvittie} in the early 1960s and revisited in the 1980s and 1990s; see the textbook summary in \\citet{fpoc}.  Just as redshift is direct, kinematical evidence for cosmic expansion, redshift drift is likewise for cosmic acceleration.  This directness is an attractive feature as it does not depends on dynamics, i.e.\\ the equations of motion (and separation of matter and dark energy density). Therefore, even if such a cosmic probe happens not to reach practically the same accuracy on dynamical cosmological-model parameters as more established distance or growth of structure probes,  it is still worthwhile exploring possibilities for carrying it out.  Conventionally, this is thought of in a brute force approach: stare at an object for a long time and measure the change in its redshift. Since the time scale for an order unity variation is the characteristic time scale of the expansion, the Hubble time of $\\sim10^{10}$ y, this requires fantastically accurate measurements of redshift stable over the observing period.  Moreover, any other time variation in the metric that is not linear gives a competing effect.  From the Principle of Equivalence, any acceleration of the source or change in gravity along the line of sight acts as a systematic contribution.  Here we explore two parallel tracks to make this cosmological probe more viable. One is exploring theoretical ideas for letting the universe do the difficult work for us, by taking advantage of source redshifts delivered to us at effectively different epochs and of differential rather than absolute measurements.  The second involves experimental approaches to improve the brute force precision, while again using differential measurements to ameliorate instrumental systematics. As with any probe, systematics are a key concern so any method that may enable better control of them is important to consider.  In Sec.~\\ref{sec:basic} we review the basic redshift drift and the plethora of systematics that confront it, as well as the potential leverage on cosmological model parameters.  Section~\\ref{sec:methods} outlines three theoretical alternatives to the brute force approach, each with their own advantages and speculative aspects. We turn to some innovative experimental approaches in Sec.~\\ref{sec:tech}, which may offer improvements in precision and systematics, and conclude in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Direct kinematic measurements of cosmic acceleration can provide valuable, model independent confirmation of dark energy. The accuracy required for detecting redshift drift is challenging, and we have reviewed some of the systematics, presented a series of theoretical speculations for observational probes, and explored various promising improvements to experimental approaches to measuring directly spectral line shifts with high accuracy.  One result is that the greatest leverage on dark energy properties comes from low-redshift observations, due to lesser covariances there with other cosmological parameters. The dark energy figure of merit can exceed 100 for a set of five measurements over $z=[0.1-0.5]$, each at 1\\% precision (together with current precision on the Hubble constant), i.e.\\ determining $w$ to 0.02, $w_0$ to 0.06, $w_a$ to 0.43. Moreover, redshift drift has extraordinary complementarity with other probes, having the opposite degeneracy in the $w_0$-$w_a$ plane from all standard probes. For example, combination of the low-redshift drift with Planck CMB data increases the figure of merit by a factor 12, to over 1000. Even a more practical 5\\% precision on redshift drift gives a FOM of 290 when combined with current Hubble constant and CMB data. Lesser complementarity can be achieved by combining redshift drift at low and high redshifts.  Given the potential leverage of redshift drift, and the ability for low redshift measurements to be powerful, we reconsider the ``stare and wait'' method of observing spectral lines from a high redshift source (or a series of sources in the Lyman-$\\alpha$ forest) over a time span of decades. We speculate about three different concepts. The first measures Hubble drift rather than redshift drift per se, gives a complementary measure of acceleration, and can be accomplished with normal galaxy redshift survey data using differential radial baryon acoustic oscillations (drBAO). The second would use cosmic ``pulsers'', not yet discovered, but possibly realizable through binary gravitational wave sources. Finally, long time-delay strong gravitational lenses, possibly to be detected by future deep time domain surveys, are also considered. We emphasize that these ideas are speculative and need considerable development, but do have considerable synergy with surveys already planned such as DESI, Euclid, and LSST.  They can also replace patience -- of several decade time scales, see the categorization of \\citet{stebbins} -- with ideas on how to get the universe and astrophysics to work for us.  More practically, we turn to experimental approaches for improving high precision redshift measurements. Given the leverage of low redshift we focus on bright emission line galaxies as targets. We explore three alternatives to conventional spectrographs, involving differential and interferometric methods, and calculate their signal-to-noise and potential redshift precisions. We find that single source redshift measurement precision of $10^{-8}$ is potentially achievable, with further improvement possible from multiple sources, multinight exposures, and larger telescopes (our numbers are canonically for 8 hours on a 10-m telescope). ``Gold'' targets, with brighter and narrower doublet emission than the galaxies considered in this article, may yet be discovered in upcoming spectroscopic surveys and provide a speedier path toward detecting redshift drift. Spatially resolving the galaxy, say with an integral field unit, may isolate subregions with finer emission than that of the combined whole \\citep{2014arXiv1408.5765G}. Instruments such as Externally Dispersed Interferometers and External Dispersion Spatial Heterodyne Spectroscopes can avoid some of the systematics in conventional spectrographs. Testing the stability and/or calibration precision possible on these instruments is of interest to see whether they can satisfy the requirements imposed by the science.  If astrophysical systematics are high, interesting science may still come from peculiar acceleration maps, e.g.\\ of the local Universe. Other targets, such as narrow radio lines or sequences of lines from masers and molecular emission, should be explored, as should other uses of high-precision wavelength measurements, such as for atomic line catalogs \\citep{2008PhLB..660...81A, 2008MNRAS.391.1308Q, 2012PhR...521...95Q}.  Cosmic acceleration is a fundamental mystery of physics, and the possibility of adding a new probe of it -- one that is kinematic and does not depend on assumptions of how to separate matter and dark energy, or of their evolution -- is an exciting prospect. The hardware and systematics will be highly challenging, but success means that we will open up a new time domain: the ability to watch the universe as it evolves in real time."
    },
    "1402/1402.0204_arXiv.txt": {
        "abstract": "At intermediate and high degree $l$, solar p- and f modes can be considered as surface waves. Using variational principle, we derive an integral expression for the group velocities of the surface waves in terms of adiabatic eigenfunctions of normal modes, and address the benefits of using group-velocity measurements as a supplementary diagnostic tool in solar seismology. The principal advantage of using group velocities, when compared with direct analysis of the oscillation frequencies, comes from their smaller sensitivity to the uncertainties in the near-photospheric layers. We address some numerical examples where group velocities are used to reveal inconsistencies between the solar models and the seismic data. Further, we implement the group-velocity measurements to the calibration of the specific entropy, helium abundance Y and heavy-element abundance Z in the adiabatically-stratified part of the solar convective envelope, using different recent versions of the equation of state. The results are in close agreement with our earlier measurements based on more sophisticated analysis of the solar oscillation frequencies (Vorontsov et al. 2013, MNRAS 430, 1636). These results bring further support to the downward revision of the solar heavy-element abundances in recent spectroscopic measurements. ",
        "introduction": "The major difficulty in the seismic measurements of the solar internal structure comes from the uncertain effects of the outermost solar layers (the photosphere and layers immediately below), where trapped acoustic waves are reflected to the solar interior. The difficulty originates from both the uncertainties in the theoretical modeling of these layers (e. g. effects of the penetrative convection), and from poor understanding of the physics of wave propagation there (non-adiabatic effects). This difficulty is behind the dominant source of mismatch between the observational and theoretical p-mode frequencies. The discrepancy increases with frequency (as the upper turning points of the p modes move upwards), reaching values of the order of one percent at frequencies of maximum oscillation power (about 3 mHz). In standards of solar seismology, one percent is a huge quantity: the frequencies of solar oscillations are measured with precision better than one part in $10^4$, and it is this high precision which enables the seismic data with its unique diagnostic capability.  To allow an accurate diagnostic of the deep interior, the near-surface uncertainties are suppressed, in one way or another, by a proper design of helioseismic inversion technique. Principally, the separation of the uncertain effects is made possible by the relatively small values of the sound speed in the subsurface layers, which makes the acoustic ray paths nearly vertical there, when the degree $l$ of the oscillations is not too high. The possibility of separating the uncertain effects is most transparent when high-frequency asymptotic analysis is implemented to describe the solar p modes: the subsurface effects bring a frequency-dependent phase shift $\\alpha(\\omega)$ of the standing acoustic wave, which does not depend on the degree $l$ when $l$ is small.  Separation of the near-surface uncertainties comes for a price of loosing valuable diagnostic information. An example is He II ionization region, the domain which is particularly important for measuring the solar abundances and for the calibration of the equation of state. For p modes of low degree $l$, the signal of He ionization is also seen as a frequency-dependent ``surface phase shift''; this signal is suppressed together with near-surface uncertainties when an arbitrary function of frequency is allowed for $\\alpha(\\omega)$. At higher degree $l$, we meet another difficulty. In high-precision measurements, the acoustic waves in the subsurface layers can no longer be considered as purely vertical, and at least a first-order correction shall be added to $\\alpha(\\omega)$ to account for the resulted dependence of the surface phase shift on the degree $l$, as another function of frequency multiplied by $l(l+1)$ \\citep{Brodsky93}. But according to the asymptotic description, allowing the degree dependence to the surface term widens the family of possible solutions in the deep interior \\citep{Gough95}.  It is desirable, therefore, to extend the set of diagnostic tools, implemented in solar seismology for analyzing the oscillation frequencies, by adding new tools which respond differently to the near-surface uncertainties, and suppress these uncertainties in a different way. This will allow more options for the cross-validation of the results, to make them more reliable. An issue of particular importance is possible effects of systematic errors in frequency measurements; these errors may be significantly bigger than the reported observational uncertainties \\citep[see e. g.][which we refer below as Paper I]{Vorontsov13a}. In general, systematic errors propagate differently to the results when different techniques of data analysis are implemented, and using different tools brings better chances to detect these errors.  In this paper, we consider solar oscillations of intermediate and high degree $l$ as surface waves, and address the diagnostic properties of group velocities of these waves. The concept of group velocity is known to be a valuable tool in terrestrial seismology, where it is applied to study the propagation of Love's and Rayleigh's waves \\citep[see e. g.][]{Dahlen98}. Section 2 contains a general discussion, based on the integral representation of the group velocity in terms of adiabatic eigenfunctions of normal modes, developed in the Appendix. In section 3, we test the diagnostic potential of group velocities by addressing the agreement of several solar models with observational data. In section 4, we implement the group-velocity analysis to the calibration of the main parameters of the solar convective envelope: specific entropy in the adiabatically-stratified layers, helium abundance $Y$ and heavy-element abundance $Z$. Section 5 contains a short discussion.  \\section[]{The group velocity} With temporal dependence separated as $\\exp(-i\\omega t)$, the displacement field of the oscillations specified by a particular spherical harmonic is \\begin{equation} {\\bf u}=\\hat{\\bf r}U(r)Y_{lm}(\\theta,\\phi)+V(r)\\nabla_1Y_{lm}(\\theta,\\phi), \\end{equation} where $\\nabla_1$ is horizontal component of the gradient operator, $\\nabla_1=\\hat{\\bf\\theta}\\partial/\\partial\\theta+\\hat{\\bf\\phi}\\sin^{-1}\\theta\\,\\partial/\\partial\\phi$, and unit vectors are designated by hats. The horizontal wavenumber at the photospheric level is $k_H=L/R_\\odot$, with $L^2=l(l+1)$. The horizontal phase velocity $v_p$ and group velocity $v_g$ are \\begin{equation} v_p={\\omega\\over L}R_\\odot,\\quad\\quad v_g=\\left({\\partial\\omega\\over\\partial L}\\right)_nR_\\odot, \\end{equation} where the derivative is taken at constant radial order $n$.  Using self-adjoint properties of the equations of linear adiabatic oscillations, it is shown in the Appendix that \\begin{equation} {v_g\\over v_p}\\equiv \\left({\\partial\\ln\\omega\\over\\partial\\ln L}\\right)_n= {\\int\\limits_0^R\\rho_0r^2L^2V^2dr +{L^2\\over 4\\pi G\\omega^2}\\int\\limits_0^\\infty P^2dr \\over \\int\\limits_0^R\\rho_0r^2\\left(U^2+L^2V^2\\right)\\,dr}\\,, \\end{equation} where $P=P(r)$ describes the Eulirean perturbation $\\psi'$ to the gravitational potential as \\begin{equation} \\psi'=-P(r)Y_{lm}(\\theta,\\phi). \\end{equation} At high degree $l$, the effects of gravity perturbation are small, and the second term in the nominator of the equation (3) can be neglected. The first term in the nominator is proportional to the mean kinetic energy of the horizontal motions; the denominator is proportional to the total kinetic energy. When the effects of gravity perturbations are small, the ratio $v_g/v_p$ is thus the ratio of the horizontal kinetic energy to the total kinetic energy.  An important property of the group velocity (when compared with phase velocity) is its enhanced sensitivity to the stratification of the inner part of the acoustic cavity, where the horizontal kinetic energy is localized (Fig. 1). Closer to the surface, the acoustic ray paths become nearly vertical, and the kinetic energy is dominated by the vertical motion. \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig1.eps} \\end{center} \\caption{Horizontal and vertical kinetic energy densities for $p_4$ mode of $l=300$. The energy per unit depth in the integrals in equation (3) was multiplied with adiabatic sound speed $c(r)$ to account for the rescaling of the independent variable from geometrical to acoustic radius.} \\label{f1} \\end{figure}  The concept of group velocity is not restricted to high-degree modes; it extends formally to all the non-radial modes when we consider the degree $l$ as a continuous parameter. The ratio $v_g/v_p$ in the degree range $0\\le l\\le 300$, calculated for the reference solar model S of \\citet{Christensen96}, is shown in Fig.2. \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig2.eps} \\end{center} \\caption{Ratio of group and phase velocities calculated for the solar model in the degree range $0\\le l\\le 300$ and frequency range 1mHz$\\le\\omega/2\\pi\\le$5mHz. Red circles show modes with frequencies below 2mHz, green---between 2 and 3mHz, and blue---above 3mHz. Upper scale is the position of the inner turning point in radius.} \\label{f2} \\end{figure} For modes penetrating deep into the solar interior, the ratio $v_g/v_p$ tends to collapse to a single function of the penetration depth (Fig. 2), the behaviour which reflects the high-frequency asymptotic properties of solar p modes. The rapid variation seen at $r_1\\approx 0.7R_\\odot$ comes from the rapid change in the sound-speed gradient at the base of the convection zone. The prominent fluctuations in $c_g/c_p$ exhibited by modes with turning points closer to the surface are produced by the rapid variation of the adiabatic exponent in the He ionization region.  For f modes, the ratio $v_g/v_p$ is close to 1/2, which reflects equipartition of the kinetic energy between horizontal and vertical motions. Indeed, using an approximate dispersion relation of high-degree solar f modes as $\\omega^2\\simeq k_Hg_0(R_\\odot)$, where $g_0(R_\\odot)$ is surface gravity, we have $(\\partial\\ln\\omega/\\partial\\ln L)_{n=0}\\simeq 1/2$.  Introducing the derivative $(\\partial\\omega/\\partial L)_n$, we extend the family of solutions to the oscillation equations from integer to continuous values of degree $l$. This extension can be achieved simply by allowing $l$ to take arbitrary values in the ordinary differential equations resulted from variable separation; it can be viewed as a result of relaxing the periodic boundary conditions in angular coordinates.  We now proceed with similar generalization but allowing non-integer values to the radial order $n$. This is equivalent to allowing a continuous variation to the radial phase integral, which takes values of $\\pi(n+1)$ at resonant frequencies. The only limitation here is that the phase has a well-defined measure only for functions with nearly-harmonic behaviour. Such a behaviour is exhibited by the solutions to the oscillation equations in the adiabatically-stratified part of the solar convective envelope, owing to their high-frequency asymptotic properties (the wave propagation here is close to that of purely acoustic waves). Specifically, the near-harmonic behavior is exhibited with best accuracy by eigenfunction $\\psi_p(\\tau)$ defined as \\citep[see][and Paper I]{Vorontsov91b} \\begin{equation} \\psi_p=\\rho_0^{-1/2}r\\left({1\\over c^2}-{\\tilde w^2\\over r^2}\\right)^{1/4} \\left(1-{N^2\\over\\omega^2}\\right)^{-1/2}p_1, \\end{equation} where $p_1=p_1(r)$ describes the Eulerian pressure perturbations $p'$ as \\begin{equation} p'=p_1(r)Y_{lm}(\\theta,\\phi), \\end{equation} \\begin{equation} \\tilde w^2={L^2\\over\\omega^2}, \\end{equation} (this parameter specifies the radial position $r_1$ of the inner turning points), and independent variable $\\tau$ satisfies \\begin{equation} {d\\tau\\over dr}={1\\over c}\\left(1-\\tilde w^2{c^2\\over r^2}\\right)^{1/2}. \\end{equation} Two linearly-independent solutions to the oscillation equations in Cowling approximation (which is applicable in the low-density envelope) are $\\psi\\simeq\\sin(\\omega\\tau)$ and $\\psi\\simeq\\cos(\\omega\\tau)$.  Using variational principle for evaluating the variation with frequency of phases of the inner and of the outer solutions (which satisfy inner and outer boundary conditions, respectively) at a boundary taken in the domain where asymptotic description is applicable, we obtain (see Appendix) \\begin{equation} \\left({\\partial\\omega\\over\\partial n}\\right)_L ={\\pi\\over 2\\omega^2}{\\psi_p^2+{1\\over\\omega^2}\\left({d\\psi_p\\over d\\tau}\\right)^2 \\over\\int\\limits_0^R\\rho_0r^2\\left(U^2+L^2V^2\\right)dr}. \\end{equation} In analogy with $(\\partial\\omega/\\partial L)_n$, considered as an angular component of the group velocity, $(\\partial\\omega/\\partial n)_L$ can be considered as a mean group velocity in radial direction.  The dependence of both the $(\\partial\\omega/\\partial L)_n$ and $(\\partial\\omega/\\partial n)_L$ on the near-surface uncertainties comes principally from the dependence on these uncertainties of the denominator in the expressions (3) and (9), which is the same (mode's kinetic energy). This observation suggests using the ratio \\begin{equation} \\gamma(L,n)=\\left({\\partial\\omega\\over\\partial L}\\right)_n\\Big/\\left({\\partial\\omega\\over\\partial n}\\right)_L \\end{equation} as diagnostic quantity, which will allow to suppress the effects of the near-surface uncertainties. In the simplest way, this quantity can be evaluated from the mode frequencies using central differences as \\begin{equation} \\gamma_{ln}={\\omega_{l+1,n}-\\omega_{l-1,n}\\over\\omega_{l,n+1}-\\omega_{l,n-1}}. \\end{equation} According to the equation (9), the inverse of $(\\partial\\omega/\\partial n)_L$ can be considered as ``mode mass''. The difference with traditional definition of the mode mass (mode energy at unit surface amplitude) is that the surface amplitude is replaced by the amplitude in the propagation domain. The principal advantage is that with this definition, the ``mode mass'' becomes an observable quantity.  The diagnostic properties of $\\gamma(L,n)$ defined by the equation (10) can be seen better if we extend the comparison with high-frequency asymptotic analysis somewhat further. In the leading-order approximation, the asymptotic eigenfrequency equation is \\begin{equation} \\omega F(\\tilde w)=\\pi\\left[n+\\alpha(\\omega)\\right], \\end{equation} where \\begin{equation} F(\\tilde w)=\\int\\limits_{r_1}^R\\left(1-\\tilde w^2{c^2\\over r^2}\\right)^{1/2}{dr\\over c}. \\end{equation} The left-hand side of the equation (12) is the radial phase integral $\\int_{r_1}^Rk_rdr$ of a purely acoustic wave, $\\alpha(\\omega)$ is the frequency-dependent ``surface phase shift''. The sound-speed profile $c(r)$ can be recovered from $dF(\\tilde w)/d\\tilde w$ using Abel's integral transform applied to the equation (13).  Differentiating both sides of the equation (12) in frequency $\\omega$, first at $L$=const, then at $n$=const, and subtracting the results, we have \\begin{equation} {dF\\over d\\tilde w}=-\\pi\\left({\\partial\\omega\\over\\partial L}\\right)_n\\Big/\\left({\\partial\\omega\\over\\partial n}\\right)_L, \\end{equation} and we see that the influence of the unknown behaviour of $\\alpha(\\omega)$ on the results of the sound-speed inversion is successfully eliminated \\citep[it is indeed the way in which some first helioseismic sound-speed inversions were performed, see e. g.] []{Vorontsov89}. Comparing the equations (10) and (14), we see that $\\gamma(L,n)$ is expected to be largely insensitive to the near-surface uncertainties.  \\section[]{Testing solar models with seismic data} In this section, we test the ability of group-velocity measurements to reveal inconsistencies between the solar models and the observational data. In these tests, we compare the $\\gamma_{ln}$-values measured from the solar oscillation frequencies with those obtained from the eigenfrequencies of solar models.  Fig. 3(a) shows the difference in $\\gamma_{ln}$ between the Sun and the reference model S of \\citet{Christensen96}. The observational frequencies were obtained by averaging the results of 15 years of SOHO MDI measurements (this observational data set is discussed in more detail in Paper I, where it is designated as data set 1). A prominent mismatch is seen in the group-velocity data for waves with turning points just below the convective envelope, which indicates an inadequate description of the seismic stratification in the solar tachocline. \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig3.eps} \\end{center} \\caption{(a) difference between observational values of $\\gamma_{ln}$ and those of the reference model S. (b) same residuals but obtained with interpolated values of $\\gamma_{ln}$ (see text) (c) as (b), but after a slowly-varying function of frequency was subtracted from the residuals. Red circles show the results obtained at frequencies below 2 mHz, green circles---between 2 and 3 mHz, and blue circles---with data above 3 mHz.} \\label{f3} \\end{figure}  For waves confined in the convective envelope, the mismatch is moderately small at low frequencies (below 2 mHz), but grows significantly when frequency increases. This behaviour is induced by the systematic difference between observational and theoretical frequencies, which grows when frequency increases. As a result, the $\\gamma_{ln}$-values evaluated with using the equation (11) from observational and theoretical frequencies correspond to surface waves with different penetration depth. A simple way to reduce this effect is to replace the theoretical values of $\\gamma_{ln}$ with values obtained by the interpolation along the p-mode ridge to proper values of $\\tilde w$ (defined by equation 6). The result is shown in Fig. 3(b); as expected, the residuals now fall much closer to a single function of the penetration depth. Small, but systematic fluctuations around the common trend remain in the residuals even after the interpolation. These are due to the fact that that the observational and theoretical values of $\\gamma$, reduced to the same penetration depth, are still measured at different frequencies (as a result, the phase of the wave function in the He II ionization region is distorted). We can make the signal which brings information about inconsistencies in deep interior cleaner still by subtracting a common function of frequency: the result is shown in Fig. 3(c) (in this computation, $f(\\omega)$ was obtained by approximating the residuals in the domain $r_1>0.85 R_\\odot$ by polynomial of degree 10 in frequency $\\omega$).  The model was than corrected by helioseismic inversion to bring it into agreement with seismic data. The resulted correction to the sound-speed profile is shown in Fig. 4(a). The inversion technique is described in Paper I; it results in another hydrostatic model, which allows a new set of eigenfrequencies to be calculated. We processed these frequencies in the same way as discussed above, to test the new model against the solar data. This is an important test: the solution in the deep interior may not be unique because two arbitrary functions of frequency were allowed by the inversion to account for near-surface effects. The result is shown in Fig. 4(b): there is no signal in the residuals. As the solar data were processed by the inversion in a very different manner, we now have better confidence in the results. \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig4.eps} \\end{center} \\caption{(a) difference in the sound speed between the Sun and the reference model S, obtained in helioseismic structural inversion with observational frequencies. (b) As Fig. 3(c), but obtained with eigenfrequencies of the model resulted from the inversion.} \\label{f4} \\end{figure}  Fig. 5 shows the results of similar tests but performed with two solar models having nearly-optimal parameters of the adiabatically-stratified part of the convective envelope (specific entropy and chemical composition) which were measured by the seismic calibration described in Paper I.  Note that the vertical scale differs by an order of magnitude from that used in Fig. 3. A small but significant mismatch with observations is seen in both the two models. Comparing with Fig. 3(c) and with corresponding sound-speed difference (Fig. 4a), we can say that the solar sound speed is slightly bigger than in the first model (Fig. 5a) in the domain $0.85 R_\\odot<r<0.9 R_\\odot$ and slightly smaller in the domain $0.8 R_\\odot<r<0.85 R_\\odot$ (we can say nothing about deeper layers because, as with model S, the residuals become distorted by much bigger inaccuracy in the tachocline). The second model (Fig. 5b) shows better agreement in the interval $0.8 R_\\odot<r<0.85 R_\\odot$, but disagreement in the interval $0.85 R_\\odot<r<0.9 R_\\odot$ is made bigger. These two models were used as a reference in the structural inversions for the adiabatic exponent $\\Gamma_1$ in Paper I. The results are shown in Paper I by Figs 12(a) and 12(b) for the first and the second model, respectively; they are in agreement with the conclusions drawn from the analysis of group velocities. The magnitude of the mismatch between the models and the data is quite small: it calls for corrections of about $1\\cdot 10^{-4}$ in the profile of the adiabatic exponent. \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig5.eps} \\end{center} \\caption{As Fig. 3(c), but obtained with eigenfrequencies of two solar models with nearly-optimal parameters of the adiabatic part of the convective envelope, constructed with SAHA-S3 equation of state and discussed in Paper I. (a): model with $Y=0.250,\\,\\,Z=0.008$; (b) model with $Y=0.245,\\,\\,Z=0.010$.} \\label{f5} \\end{figure}  \\section[]{Calibration of the envelope model} We now implement the group-velocity measurements to the calibration of the the global parameters of the solar convective envelope---specific entropy in the adiabatically-stratified layers and two parameters of the chemical composition. As in Paper I, we compare with seismic data the 3-dimensional grids of envelope models calculated with four different versions of the equation of state. The models are described in detail in Paper I.  Fig. 6 illustrates the potential possibility of simultaneous measurement of the parameters of the convective envelope by showing the response of the mismatch in $\\gamma_{ln}$ between the Sun and the model to the variation of the specific entropy, helium abundance, and heavy-element abundance in the model. Fig. 6(a) shows the residuals obtained with the best-fit model in the grid of models calculated with SAHA-S3 equation of state. Figs 6(b) and 6(c) show the response of the residuals to the variation of the specific entropy (controlled by a mixing-length parameter $\\alpha$) and to the variation of the helium abundance $Y$. Each of the two variations produces a quasi-periodic signal in the residuals by changing the profile of the adiabatic exponent in the He II ionization region. The two signals look similar, but closer inspection reveals that they have different phase: smaller entropy (bigger $\\alpha$) shifts the He II ionization to greater depths (it also shifts the location of the upper turning points, but these two effects do not compensate for each other). \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\linewidth]{fig6.eps} \\end{center} \\caption{Difference between solar values of $\\gamma_{ln}$ and model predictions, for the ``best-fit'' model in the 3-D grid of envelope models (a), for a model of the same chemical composition but with different specific entropy in the adiabatically-stratified layers (b), model which differs from the best-fit model in He abundance Y (c), and model which differs in the heavy-element abundance Z (c). Red circles show the results in the $\\tilde w$-range between 4000 s and 7000 s (lower turning points between $0.85 R_\\odot$ and $0.933 R_\\odot$), blue circles---for $\\tilde w>7000$ s.} \\label{f6} \\end{figure} Principally, it is the availability of both the amplitude and phase of the He II ionization signal in the solar oscillation frequencies which allows separate measurement of helium abundance and entropy; this property was used in the first seismic measurements of the solar He abundance \\citep[e. g.][]{Vorontsov91} \\citep[for extended discussion, see][]{Vorontsov92}. The signal produced by the variation of the heavy-element abundance (Fig. 6d) is more complicated; a distinctive feature of this signal is that bigger $Z$ shifts $\\Delta\\gamma_{ln}$ to negative values, which signals that the sound speed in the model is too small (smaller $Z$ brings smaller values to the adiabatic exponent $\\Gamma_1$).  The results of the calibration are illustrated by Fig. 7. In these computations, the frequency range of the input data was limited by lower frequencies (below 2 mHz), where the difference in absolute values of the observational and theoretical frequencies is relatively small, and observational and theoretical values of $\\gamma_{ln}$ were compared directly, without any interpolation to common penetration depths (common values of $\\tilde w$). The goodness of fit was measured by the merit function ($\\chi^2$ per degree of freedom) \\begin{equation} M^2={1\\over N}\\sum\\limits_{l,n} \\left[{\\gamma_{ln}^{\\rm(obs)}-\\gamma_{ln}^{\\rm(model)}\\over\\delta\\gamma_{ln}^{\\rm obs}}\\right]^2, \\end{equation} where $\\delta\\gamma_{ln}^{\\rm obs}$ is the $1\\sigma$ uncertainty in the observational values of $\\gamma_{ln}$ (defined by equation 11) induced by the expected random errors in the oscillation frequencies. The solar p-mode frequencies were measured from the 1-yr SOHO MDI power spectra by the technique described in \\citep[][]{Vorontsov09,Vorontsov13b} (data set 4 of Paper I). As in Paper I, the input date were limited by $\\tilde w>4000$ s (inner turning points $r_1>0.85 R_\\odot$) to eliminate modes with theoretical frequencies distorted by the inadequate description of the solar tachocline. (The same data were used in the results shown in Fig. 6). \\begin{figure} \\begin{center} \\includegraphics[width=0.95\\linewidth]{fig7.eps} \\end{center} \\caption{Goodness of fit (merit function $M$) of envelope models calculated with (a) OPAL-1996, (b) OPAL-5005, (c) SAHA-S2 and (d) SAHA-S3 equations of state. $M_{\\rm opt}$ is the best value of the merit function, $\\Delta M$ is the interval between contour lines (solid curves). For each pair of $Y$ and $Z$, the specific entropy was choosen to optimize $M$. Dashed level lines show the mass coordinate $m_{0.75}$ (equation 16) in the models optimized with respect to the specific entropy. The thick dashed line is for $m_{0.75}=0.9822$.} \\label{f7} \\end{figure}  Dashed lines in Fig. 7 show the values of the dimensionless mass coordinate taken at $r=0.75 R_\\odot$, \\begin{equation} m_{0.75}=m(0.75 R_\\odot)/M_\\odot, \\end{equation} for models which have optimal specific entropy. When helioseismic structural inversion is performed into the radiative interior, this parameter determines the density profile obtained in the solar core. The ability of the inversion to fit low-degree measurements depends on a proper value of this parameter (corresponding effects in the oscillation frequencies come from the effects of gravity perturbation in the high-density core). Successful inversion into the deep interior requires $m_{0.75}\\simeq 0.9822$ (this finding is discussed in more detail in Paper I).  The maximum-likelihood values of $Y$ and $Z$, resulted from the calibration, are in agreement with those reported in Paper I; they are in the range of $Y$=0.245--0.260 and $Z$=0.006--0.011. On average, the optimal values of $Y$ are now slightly bigger (by about 0.005), and optimal values of $Z$ are slightly smaller (by about 0.001).  The major difference with the results reported in Paper I is that different versions of the equation of state allow to achieve nearly the same optimal values for the merit function. In our vision, the failure of our group-velocity analysis to distinguish between the performance of different equations of state comes principally from the more limited amount of the input data. The entire (and most valuable) p$_1$-ridge, for example, is only used for evaluating the group velocities of $n=2$ waves (see equation 11). Also, evaluating $(\\partial\\omega/\\partial n)_L$ using central differences over large frequency intervals (up to 1 mHz) brings an excessive averaging and loss of spatial resolution.  The attractive feature of the calibration with $\\gamma_{ln}$ is its ultimate simplicity, which brings better confidence to the results. Two functions of frequency were allowed in the calibration of Paper I to account for the uncertain near-surface effects; in general, this strategy makes the results more ambiguous. No allowance for any uncertainties was given in the calibration which is described above. Another convenient feature is that $\\gamma_{ln}$, as dimensionless quantity, is invariant to the homology rescaling of the hydrostatic model; the calibration reported in Paper I had to implement the rescaling as an extra (fourth) parameter, to allow small corrections to the solar radius.  Since the near-surface uncertainties are not eliminated completely in the analysis, we performed several numerical experiments to address the stability of the calibration. Fig. 8(a) shows shows the result obtained with using interpolated model values of $\\gamma_{ln}$ (see section 3). Fig. 8(b) shows the result obtained when the grid of models was recalculated using \\citet[][]{Canuto91} convection theory as an alternative to standard prescription. This modification brings the absolute values of the theoretical frequencies to somewhat better agreement with observations. In Fig. 8(c) we address the result of an artificial experiment where outer boundary conditions in the eigenfrequency computations (here, we implemented ``zero'' boundary conditions, with setting Lagrangian pressure perturbation to zero) were shifted from the temperature minimum to the photospheric level. This modification makes the discrepancy between observational and theoretical frequencies bigger. Fig. 7(d) shows the result obtained when the observational frequency set was replaced with earlier measurements \\citep[][data set 3 of Paper I]{Schou99}. We conclude from these experiments that calibration is moderately stable, but the accuracy of measuring the chemical-composition parameters of the solar envelope is not better than 0.005 in $Y$ and 0.002 in $Z$. \\begin{figure} \\begin{center} \\includegraphics[width=0.95\\linewidth]{fig8.eps} \\end{center} \\caption{Stability of the calibration of envelope models calculated with SAHA-S3 equation of state (each panel has to be compared with Fig. 7d). (a) the result obtained when observational and theoretical values of $\\gamma_{ln}$ were compared with using interpolation to common values of $\\tilde w$; (b) the result obtained with models calculated using an alternative prescription of the convection theory; (c) effect of changing the outer boundary conditions in eigenfrequency computations; (d) effect of changing the observational data set (see text).} \\label{f8} \\end{figure}  \\section[]{Discussion} The analysis of the solar p-mode data in terms of group velocities of surface waves represents an alternative tool of helioseismic measurements, which can be used productively for validating the results obtained with more traditional methods of solar seismology. The distinctive feature of the group velocities is their enhanced sensitivity to the solar stratification in the bottom part of the acoustic cavity, and smaller sensitivity to the uncertain effects of the near-photospheric layers, when compared with oscillation frequencies. This property allows to make the technique of seismic analysis more simple and transparent.  Calibration of the chemical-composition parameters $Y$ and $Z$ in the solar convective envelope confirms our previous results (Paper I), obtained with using a much more sophisticated analysis of the solar oscillation frequencies. All our results support strongly the downward revision of heavy-element abundances reported in recent spectroscopic measurements \\citep[][]{Asplund09}.  In further work, the development of inversion techniques implementing the concept of group velocity may be a significant step forward. Another issue is related with raw data analysis at high degree $l$, where accurate frequency measurements represent a very difficult task \\citep[for an account of the current efforts, see][]{korz13}. An interesting approach may consist in measuring the group velocity directly, as a slope of the p-mode ridge in the $l-\\nu$ power spectra.  \\section*{Acknowledgments} In this work, V.A.B. and S.V.A. were supported by the RBRF grant 12-02-00135-a.  \\begin{thebibliography}{99}  \\bibitem[\\protect\\citeauthoryear{Asplund et al}{2009}]{Asplund09} Asplund M., Grevesse N., Sauval A.J., Scott, P., 2009, Ann. Rev. A\\&A, 47, 481  \\bibitem[\\protect\\citeauthoryear{Brodsky \\& Vorontsov}{1993}]{Brodsky93} Brodsky M., Vorontsov S. V., 1993, ApJ, 409, 455  \\bibitem[\\protect\\citeauthoryear{Canuto \\& Mazzitelli}{1991}]{Canuto91} Canuto V. M., Mazzitelli L., 1991, ApJ, 370, 295  \\bibitem[\\protect\\citeauthoryear{Christensen-Dalsgaard et al}{1996}]{Christensen96} Christensen-Dalsgaard J. et al., 1996, Sci, 272, 1286  \\bibitem[\\protect\\citeauthoryear{Dahlen \\& Tromp}{1998}]{Dahlen98} Dahlen F. A., Tromp J., 1998, Theoretical Global Seismology, Princeton Univ. 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M., 2013, ApJ, 778:75  \\bibitem[\\protect\\citeauthoryear{Vorontsov et al}{1991}]{Vorontsov91} Vorontsov S. V., Baturin V. A., Pamyatnikh A. A., 1991, Nature, 349, 49  \\bibitem[\\protect\\citeauthoryear{Vorontsov et al}{1992}]{Vorontsov92} Vorontsov S. V., Baturin V. A., Pamyatnikh A. A., 1992, MNRAS, 257, 32  \\bibitem[\\protect\\citeauthoryear{Vorontsov et al}{2009}]{Vorontsov09} Vorontsov S. V., Jefferies S. M., Giebink C., Schou J., 2009, Astron.Soc.Pacific Conf.Series, 416, 301  \\bibitem[\\protect\\citeauthoryear{Vorontsov et al}{2013}]{Vorontsov13a} Vorontsov S. V., Baturin V. A., Ayukov S. V., Gryaznov V. K., 2013, MNRAS, 430, 1636 (Paper I)  \\end{thebibliography} \\appendix ",
        "conclusions": ""
    },
    "1402/1402.7338_arXiv.txt": {
        "abstract": "We calculate the evolution of close binary systems (CBSs) formed by a neutron star (behaving as a  radio pulsar) and a normal donor star, evolving either to helium white dwarf (HeWD) or ultra short orbital period systems. We consider X-ray irradiation feedback and evaporation due to  radio pulsar irradiation. We show that irradiation feedback leads to cyclic mass transfer episodes, allowing CBSs to be observed in-between as binary radio pulsars under conditions  in which standard, non-irradiated models predict the occurrence of a low mass X-ray binary. This behavior accounts for the existence of a family of eclipsing binary systems known as redbacks. We predict that redback companions should almost fill their Roche lobe, as observed in PSR~J1723-2837. This state is also possible for systems evolving with larger orbital periods.  Therefore, binary radio pulsars with companion star masses usually interpreted as larger than expected to produce HeWDs may also result in such {\\it quasi~-~Roche Lobe Overflow} states, rather than hosting a carbon-oxygen WD.  We found that CBSs with initial orbital periods $\\mathrm{P_{i}<1}$~day evolve into redbacks. Some of them produce low mass HeWDs, and a subgroup with shorter $\\mathrm{P_{i}}$ become black widows (BWs). Thus, BWs descent from redbacks, although not all redbacks evolve into BWs.  There is mounting observational evidence favoring that BW pulsars are very massive ($\\mathrm{\\gtrsim 2\\; M_{\\odot}}$). As they should be redback descendants, redback pulsars should also be very massive, since most of the mass is transferred before this stage. ",
        "introduction": "Close binary systems (CBSs) formed by a neutron  star (NS)  orbiting together with a normal  donor star experience a mass transfer episode when the donor star undergoes Roche Lobe OverFlow (RLOF). Since then on, stars are forced to exchange  mass. A fraction $\\mathrm{\\beta \\leq 1}$ of the transferred  matter, carrying a substantial amount of angular momentum, is accreted by the NS, which becomes recycled and may reach  spin periods in the millisecond range. In the standard treatment the initial  RLOF is a long standing episode; the donor star transfers a  large fraction of its mass and the orbit is sensibly affected. During RLOF episodes, a CBS can be detected as a low mass X-ray binary (LMXB). For CBSs with initial orbital periods $\\mathrm{P_{i} >  1}$~day, after the end of the RLOF, the pair may be observed as a radio pulsar with a binary companion, and almost no further  orbital evolution is expected.  Within this general picture,  if we restrict ourselves to treat the case of stars that do not ignite helium, binary evolution predicts a tight donor mass-orbital period\\footnote{We denote the corresponding quantities of the NS and donor star with subscripts $\\mathrm{1}$ and $\\mathrm{2}$, respectively.} $\\mathrm{P(M_{2})}$ relation (\\citealt{1999A&A...350..928T}; \\citealt{2002ApJ...565.1107P};  \\citealt{2012MNRAS.421.2206D}). Thus, for CBSs with $\\mathrm{P_{i} >  1}$~day, standard calculations that  lead to the formation of HeWDs predicts the occurrence of companions to radio pulsars only in the neighborhood of $\\mathrm{P(M_{2})}$ relation. Currently, companions to radio pulsars more massive than predicted by  the $\\mathrm{P(M_{2})}$ relation are considered as CO or HeCO~WDs (\\citealt{2000ApJ...530L..93T}; \\citealt{2002ApJ...565.1107P}).  Two separate classes of interacting binary systems with a radio pulsar have been identified, they are called redbacks and BWs. Redbacks are eclipsing binary systems with 0.1~day~$ < P <$ 1~day  and $\\mathrm{0.2\\;M_{\\odot} < M_{2} < 0.4\\;M_{\\odot}}$. BWs  have $P$ in the same range of values, but much lower companion masses ($\\mathrm{M_{2} < 0.05\\;M_{\\odot}}$).  Standard CBS models do not take into account neither evaporation of the donor star by radio pulsar  irradiation \\citep{1992MNRAS.254P..19S}, nor X-ray  irradiation feedback \\citep{2004A&A...423..281B}. During a RLOF, matter falling onto the NS produce X-ray radiation that illuminates back the donor star, giving rise to the irradiation feedback phenomenon. If the irradiated star has an outer  convective zone, its structure is sensibly affected.  \\citet{1985A&A...147..281V} studied irradiated grey atmospheres  finding  that the entropy at deep convective layers must be the same for the  irradiated and non-irradiated portions of the star. To fulfill this condition,  the irradiated surface is partially inhibited to release energy emerging from its deep interior; i.e., the effective surface becomes smaller than $\\mathrm{4 \\pi  R_{2}^{2}}$ ($\\mathrm{R_{2}}$ is the radius of the donor star). Irradiation makes the evolution to depart from predicted by standard theory. After the onset of the RLOF, the donor star relaxes to the established conditions on a thermal  (Kelvin-Helmholtz) time scale, $\\mathrm{\\tau_{KH}= G M^{2}_{2}/(R_{2}  L_{2}})$  ($\\mathrm{G}$ is the gravitational constant, and $\\mathrm{L_{2}}$ is the luminosity of the  donor star). In some cases, the structure is unable to sustain the RLOF, becoming detached. Subsequent nuclear evolution may lead the donor star to RLOF again, undergoing a quasi cyclic behavior \\citep{2004A&A...423..281B}. Thus, irradiation feedback may lead to the occurrence of a large number of short lived RLOFs instead of a long standing one. In between, the system may reveal itself as a radio pulsar with a binary companion. Notably, the evolution of several quantities  is only mildly dependent on the irradiation feedback (e.g., the orbital period).  On the other hand, evaporation of the donor star, due to radio pulsar irradiation affects the advanced evolution of CBSs that eventually reach very short orbital periods. This effect makes  the donor star to reach very low masses and makes these systems to achieve orbital periods  large enough to be compatible with observations (\\citealt{2012ApJ...753L..33B};  \\citealt{2013MNRAS.433L..11B}).  It is the aim of this work to explore the consequences of evaporation and X-ray  irradiation feedback in the evolution of these CBSs. As stated above, CBSs  undergo cyclic RLOFs. This allow us to interpret that  companions to radio pulsars more massive than predicted by the $\\mathrm{P(M_{2})}$ relation may  also be progenitors of HeWDs or ultra short orbital period systems, observationally caught between two consecutive RLOFs. This scenario provides a natural explanation for the occurrence of redback  pulsars, as well as their evolutionary link with BWs.  The reminder of the paper is as follows. In Section~\\ref{sec:numeri}  we describe our numerical evolutionary code. In Section~\\ref{sec:resultados}  we present our models, while in Section~\\ref{sec:disc_concl} we interpret and discuss these results and give  some concluding remarks. ",
        "conclusions": "\\label{sec:disc_concl}  We computed the evolution of  CBSs formed by a NS (behaving as a radio pulsar) together with a normal donor star, for a range of donor masses for which we  expect as a final product a low mass HeWD or an ultracompact system. We included X-ray irradiation feedback and evaporation due to radio pulsar irradiation.  These important ingredients allow for the occurrence of radio pulsars with companions in evolutionary stages at which donor stars are undergoing episodic mass transfers and almost filling their Roche lobe (the quasi-RLOF state), being widely different from a WD  star. This seems to fit the case of J1723-2837 \\citep{2013ApJ...776...20C} which remains otherwise hard to interpret.  We assumed that during RLOFs, the NS acts as  an X-ray point  source \\citep{1997A&AS..123..273H}. This represents the simplest configuration, but not the only possible one \\citep{2008NewAR..51..869R}. For example, the X-ray flux may come from an accretion disk. In any case, a sequence of RLOF and detachment cycles are expectable. Notice that for all the values of $\\mathrm{\\alpha_{irr}>0}$ considered here we find a  cyclic behavior for the mass transfer. Apparently, donor stars with convective envelopes are very  sensitive to irradiation feedback, making the scenario presented  here a robust one. This is so, provided ``radio ejection'' is not operative; otherwise it would inhibit accretion onto the NS and irradiation feedback too, but evaporation should occur.  Part of our models evolve to redback conditions. This is so for the whole range of initial masses and short initial orbital periods  ($\\mathrm{P<1}$~day). For all the values of $\\mathrm{\\alpha_{irr}>0}$ we find cyclic mass transfer  episodes. A fraction of them form a low mass HeWD, but  others become BWs progenitors. Therefore, our calculations indicate that all BWs should have passed through the redbacks region (although at such stage they are not necessarily classified as redbacks if eclipses are not seen). \\citet{2013ApJ...775...27C} studied this problem, based on models including evaporation. They found that only fine tuning the evaporation rate for systems with $P < 4$~h, it is possible to find BWs as redback descendants. This result is very different from the scenario described by our models.  As stated in Section~\\ref{sec:resultados}, some BWs have been reported to harbour massive NSs. In our scenario in which BWs  are redback descendants, we infer that the NSs that behave as redbacks should be massive too. This is because most of the accretion onto the NSs should have occurred previous to reaching redback conditions. We  found CBSs corresponding to redbacks with NS masses as high as 2.4~$\\mathrm{M_{\\odot}}$.  In summary, we have calculated a set of evolutionary tracks showing the relevance of irradiation/evaporation stages for interacting CBSs. Although other relevant ingredients may still be lacking, we believe there is enough evidence that a unified picture of the redback/BW evolution is closely related to the consideration of these physical effects, allowing an interpretation of present observational evidence."
    },
    "1402/1402.2037_arXiv.txt": {
        "abstract": "This Letter details a measurement of the ionization yield ($Q_y$) of 6.7 keV \\iso{40}{Ar} atoms stopping in a liquid argon detector. The $Q_y$ of 3.6--6.3 detected $e^{-}/\\mbox{keV}$, for applied electric fields in the range 240--2130 V/cm, is encouraging for the use of this detector medium to search for the signals from hypothetical dark matter particle interactions and from coherent elastic neutrino-nucleus scattering. A significant dependence of $Q_y$ on the applied electric field is observed and explained in the context of ion recombination. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0803_arXiv.txt": {
        "abstract": "We consider the high mode density reported in the $\\delta$~Scuti star HD~50844 observed by {\\it CoRoT}.  Using simulations, we find that extracting frequencies down to a given false alarm probability by means of successive prewhitening leads to a gross over-estimate of the number of frequencies in a star.  This is due to blending of the peaks in the periodogram due to the finite duration of the time series.  Prewhitening is equivalent to adding a frequency to the data which is carefully chosen to interfere destructively with a given frequency in the data.  Since the frequency extracted from a blended peak is not quite correct, the interference is not destructive with the result that many additional fictitious frequencies are added to the data.  In data with very high signal-to-noise, such as the {\\it CoRoT} data, these spurious frequencies are highly significant.  Continuous prewhitening thus causes a cascade of spurious frequencies which leads to a much larger estimate of the mode density than is actually the case.  The results reported for HD~50844 are consistent with this effect.  Direct comparison of the power in the raw periodogram in this star with that in $\\delta$~Scuti stars observed by {\\it Kepler} shows that HD~50844 has a typical mode density. ",
        "introduction": "The $\\delta$~Scuti class of variables are dwarfs or giants with spectral types between A2 and F5.  They lie on an extension of the Cepheid instability strip with periods in the range 0.02--0.3~d.  Most of the pulsational driving in these stars is by the $\\kappa$~mechanism in the He{\\sc II} partial  ionization zone.  A large number of  $\\delta$~Sct stars have been identified in photometric time series data from the {\\it Kepler} spacecraft.  A study of these stars by \\citet{Balona2011c} has revealed that even in the center of the instability strip, no more than half the stars pulsate as $\\delta$~Sct variables.  This is a surprising finding; the reason for damping of pulsations in constant stars in the instability strip is presently not known.  Prior to the {\\it MOST, CoRoT} and {\\it Kepler} space missions, it was thought that the frequency spectra in $\\delta$~Sct stars may become very dense as the detection threshold is lowered.  The reason is that modes of high degree, $l$, which are not seen from the ground because of the low amplitudes due to cancellation effects, should be easily seen at the precision level attained for these space missions \\citep{Balona1999}. Indeed, this expectation appears to have been fulfilled in {\\it CoRoT} observations of HD~50844 \\citep{Poretti2009}.  This star appears to have a very high mode density over the whole frequency range, indicating that modes of relatively high $l$ are seen.  However, in a study of {\\it Kepler} $\\delta$~Sct stars, \\citet{Balona2011c} find that, in general, the mode density is quite moderate and that HD~50844 is probably an exception.   In order to study mode density, one needs to be sure that frequencies are correctly extracted from the data.  This is not an issue for the high amplitude peaks in the periodogram, but since the number of modes is expected to increase with decreasing amplitude, one needs to estimate the probability that a given peak is due to noise (the false-alarm probability). The significance level of a frequency can be estimated easily only for the case of an equally-spaced time series, when the Lomb-Scargle criterion can be used \\citep{Scargle1982}.  When the sampled times are not uniformly spaced, the problem is intractable and can only be solved by numerical means \\citep{Frescura2008}.  Ground-based data are seldom equally spaced and the ``four-sigma'' rule is often used to estimate the significance level \\citep{Breger1993}.  This rule states that a peak is  deemed significant if its amplitude exceeds the background noise level of the periodogram by a factor of four.  This is a purely empirical rule with no statistical foundation, but does seem to be reasonable in many cases \\citep{Kuschnig1997, Koen2010}.  The {\\it CoRoT} data are uniformly spaced and it is therefore important to apply the correct statistics.  For a time series of duration, $t$, the width of a peak in the power spectrum is inversely proportional to $t$.  Thus for a time series of a given duration, there is a maximum limit on the number of peaks per frequency interval that can be measured. Since the mode density is expected to increase with decreasing amplitude, blending of peaks due to the finite frequency resolution becomes very important. Thus the number of frequencies that can be extracted, and hence the mode density, depends not only on a correct estimation of the false alarm probability, but also on effects related to the finite frequency resolution.  These effects have not been taken into account in estimates of mode density by \\citet{Poretti2009} and \\citet{Balona2011c}.   From the above considerations, it is clear that one needs to carefully reconsider the question of significance in frequency extraction for {\\it CoRoT} and {\\it Kepler} data and, in particular, the importance of frequency resolution and close frequencies.  In this paper we use simulations at various mode densities to determine the effect of mode density on frequency extraction. We also compare results obtained with the Lomb-Scargle false alarm probability (FAP) criterion with those calculated using the four-sigma rule.  It turns out that neither of these criteria are useful if the number of modes increases with decreasing amplitude.  Finally, we compare the amplitude density in the periodogram of HD\\,50844 with those of {\\it Kepler} $\\delta$~Scuti stars. ",
        "conclusions": "Using simulations, we show that it is not possible to extract frequencies reliably down to the expected noise level using successive prewhitening. The reason is due to unresolved peaks in the periodogram which become more numerous as the amplitude decreases.  Such unresolved peaks lead to a multitude of erroneous frequencies. Prewhitening by a frequency which differs from the true frequency by a significant amount leaves many frequency residuals of lower amplitude.  In the high S/N data from space missions, these spurious residual frequencies often have significant amplitudes, leading to a cascading effect of spurious frequencies.  Simulations show that frequency extraction using the Lomb-Scargle or four sigma criterion can lead to a gross over-estimate of the true number of frequencies.  Using nonlinear optimization does not solve this problem.  It appears that the extraordinary high mode density of the {\\t CoRoT} star HD\\,50844 found by \\citet{Poretti2009} is not real, but a result of this phenomenon.  Because of the effect described above, it is not possible to distinguish between frequencies present in the star and spurious low-amplitude frequencies caused by prewhitening of unresolved frequency groups.  Thus it is not possible to count the number of individual modes with any degree of certainty below a certain amplitude level.  One can, however, compare the power in the periodogram above a given amplitude for different stars.  We made such a comparison for HD\\,50844 with 357 $\\delta$~Scuti stars in the {\\it Kepler} public archive.  It turns out that HD\\,50844 is not exceptional in this regard."
    },
    "1402/1402.1109.txt": {
        "abstract": "We investigated the optical light curve of T Pyx during its 2011 outburst through compiling a database of Solar Mass Ejection Imager (SMEI) and AAVSO observations. The SMEI light curve, providing unprecedented detail covering $t$=1.5-49 days post-discovery, was divided into four phases based on the idealised nova optical light curve; the initial rise (1.5-3.3 days), the pre-maximum halt (3.3-13.3 days), the final rise (14.7-27.9 days), and the early decline (27.9 days - -). The SMEI light curve contains a strongly detected period of 1.44$\\pm$0.05 days during the pre-maximum halt phase. These oscillations resemble those found in recent TNR models arising from instabilities in the expanding envelope. No spectral variations that mirror the light curve periodicity were found however. The marked dip at $t$$\\sim$22-24 days just before light curve maximum at $t$=27.9 days may represent the same (shorter duration) phenomenon seen in other novae observed by SMEI and present in some model light curves. The spectra from the 2m Liverpool Telescope and SMARTS 1.5m telescope were obtained from $t$=0.8-80.7 and 155.1-249.9 days, covering the major phases of development. The nova was observed very early in its rise where a distinct high velocity ejection phase was evident with derived $V_{ej}$$\\sim$4000 km s$^{-1}$ initially. A marked drop at $t$=5.7 days, and then a gradual increase occurred in derived $V_{ej}$ to stabilise at $\\sim$1500 km s$^{-1}$ at the pre-maximum halt. Here we propose two different stages of mass loss, a short-lived phase occurring immediately after outburst and lasting $\\sim$6 days followed by a more steadily evolving and higher mass loss phase. The overall spectral development follows that typical of a Classical Nova and comparison with the photometric behaviour reveals consistencies with the simple evolving pseudo-photosphere model of the nova outburst. Comparing optical spectra to X-ray and radio light curves, weak [Fe X] 6375\\AA$ $ emission was marginally detected before the X-ray rise and was clearly present during the brightest phase of X-ray emission. If the onset of the X-ray phase and the start of the final decline in the optical are related to the cessation of significant mass loss, then this occurred at $t$$\\sim$90-110 days. ",
        "introduction": "Classical novae (CNe) are interacting binary systems whose outbursts are powered by a thermonuclear runaway in accreted material on the surface of a white dwarf (WD). The secondary stars in such systems fill their Roche lobe and material is transferred onto the WD primary star via an accretion disc \\citep{sta08}. Recurrent novae (RNe) are, by definition, CNe with multiple recorded outbursts and most contain evolved secondaries \\citep{dar12}. There are 10 known Galactic RNe: T Pyx, IM Nor, CI Aql, V2487 Oph, U Sco, V394 CrA, T CrB, RS Oph, V745 Sco, and V3890 Sgr, given here in order of increasing orbital period (\\citep{sch10b}, although the period of V2487 Oph is only approximately known). T Pyx ($\\alpha$=9$^{h}$4$^{m}$41$^{s}\\!.$50, $\\delta$=$-$32$^{\\circ}$22$'$47$''\\!.$5) has the shortest known orbital period ($\\sim$0$^{d}\\!.$076 or 1$^{h}\\!.$824) of any recurrent nova (\\citealp[see][for a review]{anu08}) and the only RN that is below the cataclysmic variable (CV) period gap \\citep{sch92,pat98,uth10}.  %---PUT THIS IN THESIS----According to standard evolutionary models, a CV below the period gap should be faint and have a low accretion rate ($\\dot{M}$) driven by gravitational radiation \\citep{war02,war08}. T Pyx's pre-outburst mass transfer rate turned out to be $\\gtrsim$10$^{-8}$ M$_{\\odot}$yr$^{-1}$ \\citep{pat98,sel10} which is higher than that expected ($\\sim$10$^{-10}$ M$_{\\odot}$yr$^{-1}$) in CVs with this period \\citep{hel01}.  %---PUT THIS IN THESIS----This unusually high $\\dot{M}$ is understandable if the system is a wind-driven supersoft X-ray source (SSS) which can accelerate the binary evolution \\citep{kni00}. Although the SSS has not been seen yet in quiescence, T Pyx is a weak hard source detected in an XMM-Newton spectrum \\citep{sel10}. A recent model reveals that nuclear burning WDs do not need to be bright in soft X-rays during the phase of mass accretion to reach the Chandrasekhar mass \\citep{dis10a,new13}. Thus the absence of a SSS therefore may not invalidate \\citeauthor{kni00}'s (2001) postulate.   %\\subsection{History of outbursts} T Pyx has had previous observed outbursts in 1890, 1902, 1920, 1944, 1966/1967, and now 2011, with an extensive shell of ejected gas associated with these. The photometric and spectroscopic characteristics of the first outbursts were briefly presented in \\citet{pay57} while the 1966 event was discussed post-outburst by \\citet{cat69} and \\citet{lan70}. All the outbursts have exhibited brightness fluctuations near the optical peak and the subsequent declines have been slow. The ejection velocity observed previously was $\\sim$2000 km s$^{-1}$ in 1967 \\citep{cat69,wil82} with $t_{3}$=63 days \\citep{sch10b}. The distance has been revised most recently by \\citet{sho11} to be $\\ge$4.5 kpc, with a strict lower limit of 3.5 kpc which was the previously accepted value. Taking M$_{bol}$=$-$7.0 given by \\citet{sch10b} gives $L_{bol}$ $\\sim$ 2$\\times$10$^{38}$ erg s$^{-1}$ at maximum light ($\\sim$$L_{edd}$ for a 1M$_{\\odot}$ WD).  T Pyx is unique for having an outburst light curve plateau in the decline stage which is different from the five other RNe (IM Nor, CI Aql, V2487 Oph, U Sco, and RS Oph) with observed (so-called `true') plateaus in their light curves. While the latter true plateaus are believed to be the result of the combination of the irradiation of the disc (by the supersoft emission from nuclear burning near the WD after the wind associated with the ejecta has stopped) and the steady light decline from the shell which leads to a flattening of the light curve until the nuclear burning turns off, the plateau in T Pyx is thought to arise from a different mechanism including the emission lines increasing in brightness \\citep{hac00}. Unlike true plateaus that have the flat portion starting during an apparently final decline, T Pyx's plateau starts immediately after a sharp drop ($\\sim$2.0 mag in 20 days after peak) and starts again in the final decline ($\\sim$105 days after peak). This feature, of a plateau starting after a sharp drop, is also present in RS Oph; however, the presence of a Super Soft X-ray Source (SSS) and the lack of colour changes during the plateau phase in RS Oph indicates a true plateau \\citep{hac08}. On the other hand, the shape of T Pyx's plateau in $B$ and $V$ band light curves is significantly different \\citep{sch10b}.  %---FIND A PLACE FOR THIS IN THESIS--------The time intervals between outbursts are 11.9, 18.0, 24.6, 22.1 \\citep{sch10b} and now 44.3 yr. %Another unique feature of T Pyx is the secular decline from 1890 to the present across many eruptions. It is the only RN that shows a decreasing trend in quiescent $B$ magnitude which has declined from 13.8 before the eruption in 1890 to 14.38, 14.74, 14.88, and 14.72 for the next four inter-eruption periods and finally faded to 15.5 from the 1980's until 2004 \\citep{sch05} and then to 15.7 from 2009 to 2011 \\citep{sch13}. By looking at this drop of a total of 1.9 magnitudes corresponding to a factor of 5.75 in flux, and using $\\dot{M}$$\\propto$$F^{\\beta}$ where $\\beta$=2.0 given by \\citet{sch05}, \\citet{sch10c} commented that the $\\dot{M}$ has dropped by a total factor of 33 since 1890. This convinced them to predict the next eruption to be 10$^{3}$ to 10$^{6}$ years from that time, which was of course shown to be a gross overestimate of the timescale to outburst.  %\\begin{figure}[bh!] %\\begin{center} %\\leavevmode %\\epsfxsize = 14.0cm %\\epsfysize = 14.0cm %\\includegraphics[width=30pc]{tpyx-Bmag-decline.eps} %\\end{center} %\\caption[Secular decline of $B$ magnitude and mass accretion rate] %{Secular decline of $B$ magnitude and mass accretion rate of T Pyx. The schematic plot of quiescent $B$ magnitude of previous eruptions (left) implies the mass accretion rate (right) has dropped with rate of 3$\\%$ of its earlier value (from \\citealp{sch10c}).} %\\label{tpyx-Bmag-decline} %\\end{figure}   %\\subsection{Central system and surroundings} In classical novae, the mass of the WD ($M_{WD}$) is typically around 1M$_{\\odot}$ \\citep{uth10}. For a RN, in order for the thermonuclear runaway (TNR) to occur on a short time scale, the WD has to be more massive (i.e. $M_{WD}$$\\gtrsim$1M$_{\\odot}$) and luminous with high $\\dot{M}$ \\citep{sta08a} implying that the secondary star has to be evolved. As the result of the unusually high $\\dot{M}$, T Pyx is more luminous than CNe at quiescence in which the Roche lobe filling main-sequence donor star is transferring mass onto a high-mass WD. Unfortunately, $M_{WD}$ is not accurately known for this system.  Based on optical spectroscopy obtained at the Very Large Telescope and the Magellan Telescope, \\citet{uth10} estimated a mass ratio $q$=0.2$\\pm$0.03. Thus with the mass of the secondary star $M_{2}$=0.14$\\pm$0.03M$_{\\odot}$ derived from the main-sequence mass-radius relation of a 5-Gyr isochrone given by \\citet{bar98} then the mass of the primary WD is $M_{WD}$=0.7$\\pm$0.2M$_{\\odot}$. This estimate of $M_{WD}$ is lower than that expected from previous theoretical studies which range from 1.30-1.37M$_{\\odot}$ \\citep{kat90}, 1.25-1.30M$_{\\odot}$ \\citep{sch10c}, and 1.25-1.4M$_{\\odot}$ \\citep{sel10}. Therefore \\citet{uth10} turned the problem around by mentioning that if $M_{WD}$$>$1$M_{\\odot}$ then $M_{2}$$>$0.2M$_{\\odot}$ for their estimated $q$=0.2 mass ratio.   %---FIND A PLACE FOR THIS IN THESIS--------The results of \\citet{uth10} were derived without putting any spectroscopic constraints to the mass of the secondary star since T Pyx has a bright accretion disc that outshines the spectral signature from the donor in the optical. They however used the period-density relation for a Roche lobe filling secondary star given by \\citet{egg83} to set some constraints. They adopted 0-20$\\%$ inflation to the donor star while \\citet{pat05} and \\citet{kni06} suggested the donor stars in ordinary CVs below the period gap are inflated by approximately 10$\\%$ due to mass loss relative to ordinary main-sequence stars of the same mass.   The orbital period at quiescence, $P$=0$^{d}.$076, is well established \\citep{pat98,sch92,uth10} with an increasing period $\\dot{P}=6\\times10^{-10}$ \\citep{pat98}. The light curve had a highly significant modulation in the 1966 eruption with a period of 0.099096 days found by \\citet{sch90}. Other transient periodicities at 0.1098 days and 1.24 days were found by \\citet{pat98} in quiescence. \\citet{pat98} suggested that the 1.24-day period might arise from precession in the accretion disc as it is roughly in the range of that in CVs while there is no explanation for the 0.1098-day period.  The Orbital inclination, $i$, is thought to be low from the spectral profiles and low radial velocity amplitude \\citep{uth10}. It has been estimated to be $\\sim$6$^{\\circ}$ \\citep{shah97}, 10$^{\\circ}$$-$20$^{\\circ}$ \\citep{pat98}, 20$^{\\circ}$$-$30$^{\\circ}$ \\citep{sel10}, $i$=10$^{\\circ}$$\\pm$2$^{\\circ}$ \\citep{uth10} and recently $i$=15$^{\\circ}$ \\citep{che11}. Such a fairly low inclination is suggested to be the cause of the long dip, which lasts for half of the orbit in the flat top light curve, due to the heating effect on the companion star \\citep{pat98}.  %\\begin{figure}[th!] %\\begin{center} %\\leavevmode %\\epsfxsize = 14.0cm %\\epsfysize = 14.0cm %\\includegraphics[width=30pc]{tpyx-shell-NTT-HST.eps} %\\end{center} %\\caption[T Pyx and its shell] %{T Pyx and its shell (on 18 September 1997) as seen from the ground by the ESO 5.58m NTT (left) and from space by $HST$ WFPC2 (right). Image Credit: M. M. Shara, R. E. Williams, and D. R. Zurek (STScI); R. Gilmozzi (ESO); D. Prialnik (Tel Aviv University); and NASA. } %\\label{tpyx-shell-NTT-HST} %\\end{figure}  T Pyx is the only RN that presents a persistent nova shell. This has a radius of $\\sim$5\", discovered by \\citet{due79}. A fainter outer halo around the shell discovered by \\citet{sha97} expands to a radius of $\\sim$10\". The shell has been found to be slowly expanding with thousands of discrete knots \\citep{due79,wil82,sha97} and is suggested to be a result of the shocks from the new eruption ejecta that interact with the ejecta from the previous eruption \\citep{con97}. %Using HST images over 13 years, \\citet{sch10c} inferred that a nova shell was first ejected in the year 1866$\\pm$5 at a velocity of 500-715 km s$^{-1}$ with shell total ejected mass $\\sim$10$^{-4.5}$M$_{\\odot}$. This very low velocity and too high ejected shell mass for a typical RN allowed them to conclude that the first eruption in 1866 must have been a normal classical nova eruption with a very long time of prior quiescence. T Pyx is proposed to have originally formed with a high mass WD \\citep{sch10c} and the 1866 eruption has triggered the high-$\\dot{M}$-wind-driven phase suggested by \\citet{kni00}. However this phase was not completely self-sustained \\citep{uth10} resulting in the decreasing nuclear burning rate and brightness as seen in the 1890 eruption and later.  %Figure \\ref{tpyx-shell-NTT-HST} shows T Pyx and its shell taken on 18 September 1997 as seen from the ground by the ESO 5.58m NTT (left) and from space by $HST$. The excellent angular resolution of $HST$ allow us to see that shell consists of thousands knots which are extremely slowly expanding.    %\\subsection{The outburst in 2011} The outburst of T Pyx in 2011 was discovered by AAVSO observer M. Linnolt at a visual magnitude of 13.0 on 2011 Apr 14.29 UT (JD 2455665.7931, hereafter $t$=0 day) and published in \\citet{sch13}. Figure \\ref{tpyx-lc-aavso} presents the AAVSO optical light curves of T Pyx in its 2011 outburst from before the eruption to 2012 Mar 31 (JD 2456018, $t$=353 days).  \\begin{figure}[th!] \\begin{center} \\leavevmode %\\epsfxsize = 14.0cm %\\epsfysize = 14.0cm \\includegraphics[width=30pc]{fig1} %{tpyx-lc-aavso.eps} \\end{center} \\caption[The AAVSO optical light curves of T Pyx in its 2011 outburst.] {The AAVSO optical light curves of T Pyx in its 2011 outburst.} \\label{tpyx-lc-aavso} \\end{figure}   \\citet{sch13} reported that from 2008 to 2011 Mar 31 (JD 2455652, $t$=$-$13.8 days) T Pyx's brightness was steady at $V$=15.5 with the usual periodic variation of $<$0.2 mag and it  was never brighter than $V$=15.0. On 2011 Apr 5.51 UT, it was observed with variations from $V$=14.4-14.7 and continued to fade slowly until 2011 Apr 10.54 UT (JD 2455662.0417, $t$=$-$3.7 days). There was a short brightness increase, a so called ``pre-eruption rise'' (peak around $V$=14.4 mag), occurring about 13 days prior to the initial rise \\citep{sch13}. After the sharp initial rise on day $t$=0, a small fading occurred roughly 2 days after that \\citep{sch13}.  \\citet{pat13} studied the post-outburst light curve in 2011 and found that an orbital period signal appeared by $t$=170 days ($V$=11.2) with a period increase of 0.0054(7)\\% implied to be the result of mass ejection from the WD. They derived an ejected mass of at least 3$\\times$10$^{-5}$M$_{\\odot}$, similar to that in CNe.  Spectroscopic observations of 2011 outburst are discussed in \\citet{sho11,sho13} and in \\citet{ima12}. \\citet{sho11} report high resolution observations at 7 epochs from $t$=1.6-46.6days and \\citet{sho13} describe observations at a further 4 epochs much later in the ourburst ($t$=179, 180, 349 and 360 days). \\citet{ima12} meanwhile report 11 epochs of low resolution spectroscopy at early times ($t$=2.7-30.8 days). Comparison with results of more extensive and intensive spectroscopic coverage is given below, together with high cadence photometry.  In this paper, in Section 2 we first describe our extensive photometric and spectroscopic datasets. In Section 3 we then use the results of these to explore the various phases of development of the nova both in comparison to those found in other novae and also with models of the outburst. Periodicities in the SMEI light curve are also investigated. Comparison with extensive observations at X-ray and radio wavelengths is detailed in Section 4 and in Section 5 we present our conclusions.  %Figure \\ref{tpyx-initial-rise} shows the pre-eruption rise and the initial rise. %\\begin{figure}[h!] %\\begin{center} %\\leavevmode %\\epsfxsize = 14.0cm %\\epsfysize = 14.0cm %\\includegraphics[width=30pc]{tpyx-initial-rise.eps} %\\end{center} %\\caption[The pre-eruption rise and the initial rise.] %{The pre-eruption rise (left) is believed not to be a part of the TNR. The initial rise (right) shows the TNR has started with a small break at %about 2 days later which indicates that the photosphere may have receded (from \\citet{sch13}).} %\\label{tpyx-initial-rise} %\\end{figure} ",
        "conclusions": "We investigated the optical light curve of T Pyx in its 2011 outburst through compiling a database of SMEI and AAVSO observations.  \\begin{itemize} \\item The SMEI light curve, providing unprecedented detail with high cadence data, was divided into four phases based on the idealised nova optical light curve: the initial rise; the pre-maximum halt; the final rise, and the early decline.  \\item A period of 1.44$\\pm$0.05 days was the most strongly detected and was found in the interval from the first observation to the end of the pre-maximum halt phase, before the visual maximum. We compared this result to oscillations found in CVs and ascribed to accretion disc precession \\citep{hir90}. Our observed $P_{orb}/P_{precession}$ is then 5.3$\\%$ corresponding to $q$=0.125-0.15 which compares to $q$=0.2$\\pm$0.03 derived by \\citet{uth10}. Although the period is in line with that expected from studies of disc precession in CVs, we question however whether the disc would be present so early in the outburst. Such oscillations are present in some of the light curves derived by \\citet{hil13} from TNR models and may be related to restructuring and rebalancing of the ejected envelope as it expands. We find no spectral variations related to the light curve periodicity however. \\item The pre-maximum halt and subsequent dip in the SMEI light curve at $t$$\\sim$22-24 days again resemble features in the light curves produced by \\citeauthor{hil13}, with the latter possibly mirroring the shorter duration feature seen in other novae observed by SMEI \\citep{hou10}. We note that this is coincident with a sharp transitory decline in H$\\alpha$ flux and an equally sharp increase in that from Fe II. \\end{itemize}  The spectra from the LT and SMARTS were investigated through each of the 4 phases of the optical light curve, in order to study the spectral evolution and investigate the physical causes of the variations of the light curve. We conclude, taking each phase in turn:  \\begin{itemize} \\item {\\bf Initial Rise ($t$=0.8-3.3 days).} The spectra show lines of high ionization species consistent with the presence of a high effective temperature pseudo-photosphere. The emission comes almost entirely from the continuum in this the ``fireball'' stage. The marked drop in the derived expansion velocity (4000 km s$^{-1}$ at $t$=0.8 days to $\\sim$2000 km s$^{-1}$ at $t$=2.7 days) is consistent with the initial ejection in the form of a Hubble flow, but also resembles that noted in Type Ia SNe and ascribed to interaction with pre-outburst material.  \\item {\\bf Pre-Maximum Halt ($t$=3.6-13.7 days).} The subsequent change in behaviour of the derived $V_{ej}$ during the pre-maximum halt phase may suggest 2 different stages of mass loss: a short-lived phase occurring immediately after outburst and then followed by a more steadily evolving and higher mass loss phase. The fireball spectrum is maintained until $t$=8.6 days which is also the beginning of the typical iron curtain stage. Overall, the ionization/excitation and effective temperature of the underlying pseudo-photosphere appear to be decreasing through this phase, in line with basic models.  \\item {\\bf Final Rise ($t$=14.7-27.9 days).} The typical principal spectrum of CNe seems to be apparent and displays the O I flash beginning at $t$$\\sim$17 days with the characteristics of the iron curtain phase still persisting. The gradual increase in $V_{ej}$, starting to appear in the final rise ($t$=14.8 days) and later, may be related to increasing ejection velocities from the central system, as proposed in other novae. The visual maximum at $t$=27.9 days seems to exhibit the lowest ionization lines, as expected.  \\item {\\bf Early Decline ($t$=27.9-90 days).} The Balmer lines began to have a double-peaked structure from $t$=42.7 days and this is shortly followed by the emergence of forbidden lines. The typical Orion spectrum is suggested to start at $t$$\\sim$70 days. The emission line of N V at 4603\\AA$ $ begins to emerge at $t$=73.7 days. The strong enhancement of N lines is associated with `nitrogen flaring' in the typical Orion spectrum stage.  \\item {\\bf Transition to the Nebular Phase ($t$=90-280 days).} By this time, the [O III] 5007\\AA$ $ nebular and [Fe X] 6375\\AA$ $ coronal lines have developed, the latter having been marginally detected at $t$=80 days while the rise in the X-ray light curve was detected at $t$=111 days. The last spectroscopic observation reported here at $t$=249.9 days in the nebular stage still exhibited [NII], He I, [Fe VII], and [Fe X] lines. \\end{itemize}  The overall spectral development of T Pyx is similar to that of CNe whose ejected mass is higher and velocity of ejection is lower than in typical RNe such as U Sco and RS Oph. We also found that in general the detected ionized elements are in line with those expected from the simple pseudo-photosphere models of \\citet{bat89} at the same $\\Delta$$B$ as shown in Table \\ref{bat89} and Figure \\ref{radio-x-lc}. An exception to this occurred near the end of the early decline phase where for e.g. N III and N V emission lines emission lines emerged earlier than would have been predicted.  In terms of the relationship of the optical development described here to that at other wavelengths, we found that: \\begin{itemize} \\item The rise of the high frequency (33 GHz) radio emission was detected at $t$=7-15 days during the pre-maximum halt phase in the SMEI light curve, while X-rays were not yet detected. \\item At $t$$\\sim$45 days, where the [O III] 5007\\AA$ $ was first present, the radio emission rose steeply while the X-rays were still undetectable. \\item The rise in the X-ray emission, which began at $t$=111 days and rose to peak at $t$$\\sim$144 days \\citep{kuu11}, is consistent with the appearance of lines from highly ionized species such as [Ne III] 3869\\AA$ $, [C III] 4364\\AA$ $, N III 4640\\AA$ $, found in the first blue spectrum ($t$=155.1 days) right after the seasonal gap. \\item If the onset of the X-ray phase and the start of the final decline in the optical are related to the cessation of significant mass loss, this occurred at $t$$\\sim$90-110 days. \\item During the rise in X-rays, the radio flux kept increasing and peaked at around $t$$\\sim$155 days for the highest frequency (37 GHz) corresponding to the middle of the plateau in the X-ray light curve where we clearly detected the coronal lines [Fe X] 6375\\AA$ $ and [Fe VII] 6087\\AA$ $ in our spectra ($t$=161.1 days). \\item Although the appearance of the X-ray emission is in line with predictions of the emergence of the SSS from the simple \\citet{bat89} model, we note that X-ray emission may be a mix of SSS and shocked circumstellar gas. \\end{itemize}"
    },
    "1402/1402.1290_arXiv.txt": {
        "abstract": "{% A circumstellar corona is proposed for a strange quark-cluster star during an accretion phase, that could be essential to understand the observations of the puzzling symbiotic X-ray system, 4U 1700+24. The state of cold matter at supranuclear density is still an important matter of debate, and one of the certain consequences of strange star as the nature of pulsars is the self-bound on surface which makes extremely low-mass compact objects unavoidable. In principle, both the redshifted O VIII Ly-$\\alpha$ emission line and the change of the black-body radiation area could be understood naturally if 4U 1700+24 is a low-mass quark-cluster star in case of wind-accretion. ",
        "introduction": "%  It has been more than 80 years since the first idea of stars formed of matter at nuclear (even supranuclear) density (see summaries and new achievements edited by van Leeuwen (2013) for IAU Symposium 291), but (un)fortunately, scientific challenges and opportunities still exit, especially after the discoveries of massive pulsars (Demorest et al. 2010, Antoniadis et al. 2013). Great efforts, though difficult due to unknown microphysics, are tried to model the inner structure of pulsar-like compact stars. Traditionally, quarks are believed to be confined in hadrons of neutron stars, while quark matter would exist in a core of hybrid/mixed star whose central density could be high enough to make quarks de-confined, and a radical but logical argument is that an entire compact star could be composed of quark matter, called a quark star. In addition to all of the above, among different models of pulsar structure, a quark-cluster star is a condensed object of quark clusters via residual color interaction, which distinguishes from conventional both neutron and quark stars (Xu 2010, 2013).  Hadron star and hybrid/mixed star are bound by gravity, which are covered by crusts with nuclei and electrons, whereas quark star and quark-cluster star are strongly self-bound on surface. It is then unavoidable that the mass of strange quark/quark-cluster star could be extremely low, even as low as planet-mass (so called strange planet; Xu \\& Wu 2003, Horvath 2012), although massive quark-cluster stars as high as $\\sim 3M_\\odot$ are still possible because of a very stiff equation of state of quark-cluster matter (Lai \\& Xu 2009, Lai, Gao \\& Xu 2013). Certainly, besides massive pulsars, it is also necessary to find evidence for low-mass compact objects in order to test the quark-cluster model. We would address in this paper that the discovery of the redshift (Tiengo et al. 2005, confirmed recently by Nucita et al. (2014) with publicly available X-ray data) from 4U 1700+24 could be one of such attempts.  In fact, it has been a long history to search for atomic transition from spectral features of either none-accreting (e.g., Burwitz et al. 2001) or accreting (e.g., Cottam et al. 2002) compact stars, but we have not yet seen success with certainty previously (e.g., Kong et al. 2007). The detected redshift may imply an emission cloud at a distance of $160-1000$ km from the central object if one assumes a neutron star mass of $\\sim 1.4 M_\\odot$ for 4U 1700+24 (Nucita et al. 2014). However, our solution of the system in the regime of quark-cluster star is that the central object has an extremely low mass of only $\\sim 10^{-2} M_\\odot$.  We will show that a corona-like plasma may exist above a quark-cluster star, and atomic transition lines could form in the corona if, in a high accretion phase, its density is larger than a critical value. In the limit of $z\\ll 1$, the gravitational redshift, $z$, of a compact star reads (e.g., Xu 2003), \\begin{equation} z\\simeq {GM\\over c^2R}, \\end{equation} for a stellar mass of $M$ and a radius of $R$. For low mass quark/quark-cluster stars, $M\\ll M_\\odot$, their densities are nearly uniform, so we have \\begin{equation} R=\\sqrt{{3c^2\\over 4\\pi G\\rho_\\star}z},~~M={4\\pi\\over 3}R^3\\rho_\\star,% \\label{MR} \\end{equation} where $\\rho_\\star$ is the uniform density of quark-cluster star at pressure-free state. We are applying $\\rho_\\star=2\\rho_0$ in following calculations (e.g., Lai \\& Xu 2009), with $\\rho_0$ the nuclear saturation density. For $z=0.009$, we have $R=2.3$ km and $M=0.014M_\\odot$, which hints that the compact object could be an extremely low-mass quark-cluster star, being gravity-free (see Fig. 1 of Xu 2010). ",
        "conclusions": "We propose that the compact object in the symbiotic X-ray binary, 4U 1700+24, is a low-mass quark-cluster star in a wind-accretion phase. Because of very weak interaction to convert up or down to strange quarks, most of accreted ions are bounced back after collisions with quark-clusters, and accretion material through open-field lines may finally result in a formation of secular corona above the stellar surface, with stronger ion-diffusion at higher accretion rates. Both the redshifted O VIII Ly-$\\alpha$ line and the change of the black-body emitting region could be understood naturally in this scenario.   The nature of three kinds of isolated pulsars are focused in the literatures. Their X-ray luminosity, $L_{\\rm x}$, scales from $10^{30\\sim 31}$ erg/s for the ``Magnificent Seven'', to $10^{32\\sim 33}$ erg/s for central compact objects (CCOs), and to $10^{34\\sim 35}$ erg/s for anomalous X-ray pulsars and soft $\\gamma$-ray repeaters (AXP/SGRs). The spin periods of CCOs are small, only a few hundreds of milliseconds, while that of other two are larger than a second, clustered around 10 s. CCOs are proposed to be anti-magnetar (e.g., Mereghetti 2011) in order to explain their X-ray pulsations. However, in the regime of quark-cluster stars, accretion from either inter-stellar medium or fossil disk could contribute significant heating on polar caps. Similar polar corona would also form there, with clearer pulsation for higher $L_{\\rm x}$ (note: some of CCOs may be heated also by spindown powers). In addition, the period clustering of AXP/SGRs and Magnificent Seven could also be understood with the inclusion of accretion torques (e.g., Liu et al. 2014).  It is conventionally believed that Type-I X-ray bursts are the results of unstable thermonuclear burning on the surface of neutron stars with binary accretion (e.g., Lewin et al. 1993). However, this popular point of view is challenged recently by a superburst detected in EXO 1745-248. Altamirano et al. (2012) concluded that their observations of that superburst are very challenging for current superburst ignition models even assuming a few days of low-level accretion before the superburst onset. Nevertheless, in the regime of solid quark-cluster star (e.g., Xu 2010), internal gravitational and elastic energy release during a quake of massive compact star might be responsible for triggering the bursts, manifesting as Type-I X-ray bursts. Magnetic waves (e.g., Alfv\\'en waves, Kaghashvili 1999) could effectively transport kinematic oscillation energy from the stellar interior out into the corona, heating significantly ions rather than electrons since the stellar oscillation frequency is much lower than the cyclotron frequencies of electrons and ions. Nuclear flash accompanying a starquake could also be possible. A recent study (Qu et al. 2014) of the flux evolution of three AXPs/SGRs does show that the ratios of persistent emission to the time averaged short bursts are comparable to those in the case of Type I X-ray bursts. In any case, a future extensive study on this scenario is necessary.   \\vspace{2mm}  \\noindent {\\bf Acknowledgments} This work is supported by National Basic Research Program of China (2012CB821800), NSFC (11225314) and XTP XDA04060604. An anonymous referee is sincerely acknowledged, especially for the language improving. I thank useful discussions at our pulsar group."
    },
    "1402/1402.6015_arXiv.txt": {
        "abstract": "We report on time-resolved spectroscopy of the 63 brightest bursts of \\src, detected with {\\it Fermi}/Gamma-ray Burst Monitor during its 2008-2009 intense bursting episode. We performed spectral analysis down to 4~ms time-scales, to characterize the spectral evolution of the bursts. Using a Comptonized model, we find that the peak energy, \\epeak, anti-correlates with flux, while the low-energy photon index remains constant at $\\sim-0.8$ up to a flux limit $F\\approx10^{-5}$~erg~s$^{-1}$~cm$^{-2}$. Above this flux value the \\epeak$-$flux correlation changes sign, and the index positively correlates with flux reaching $\\sim$1 at the highest fluxes. Using a two black-body model, we find that the areas and fluxes of the two emitting regions correlate positively. Further, we study here for the first time, the evolution of the temperatures and areas as a function of flux. We find that the area$-kT$ relation follows lines of constant luminosity at the lowest fluxes, $R^2\\propto kT^{-4}$, with a break at higher fluxes ($F>10^{-5.5}$~erg~s$^{-1}$~cm$^{-2}$). The area of the high$-kT$ component increases with flux while its temperature decreases, which we interpret as due to an adiabatic cooling process. The area of the low$-kT$ component, on the other hand, appears to saturate at the highest fluxes, towards $R_{\\rm max}\\approx30$~km. Assuming that crust quakes are responsible for SGR bursts and considering $R_{\\rm max}$ as the maximum radius of the emitting photon-pair plasma fireball, we relate this saturation radius to a minimum excitation radius of the magnetosphere, and put a lower limit on the internal magnetic field of \\src, $B_{\\rm int}\\gtrsim4.5\\times10^{15}$~G. ",
        "introduction": "\\label{Sec:Intro}  Soft gamma repeaters (SGRs) represent a small subset of the isolated neutron star (NS) population. They are characterized by short ($\\sim$0.1~s), bright ($10^{38}-10^{41}$~erg) bursts of hard X-ray/soft gamma-ray emission. Very rarely, they emit giant flares with extreme energies ($10^{44}-10^{46}$~erg), characterized by an initial very short, hard spike and a decaying tail lasting several minutes. Intermediate duration and luminosity bursts were also recorded for a few SGRs \\citep[e.g.,][]{kouveliotou01ApJ:IF1900}. Most SGRs are bright X-ray sources, with luminosities significantly larger than those expected from rotational energy losses. Their spin periods are clustered between 2-12~s, and their large spin-down rates imply very high surface dipole magnetic fields of $10^{14}-10^{15}$~G \\citep[e.g., ][but also see \\citealt{rea13:0418}]{kouv98Natur:1806,kouveliotou99ApJ:1900}. The above properties argue strongly in favor of the nature of SGRs as very strongly magnetized neutron stars, i.e., magnetars \\citep{duncan92ApJmagnetars, paczynski92AcA:magnetars}.  Another class of isolated NS, the Anomalous X-ray Pulsars (AXPs), show very similar characteristics to SGRs, including spin down rates leading to extreme magnetic fields. AXPs were discovered as bright persistent X-ray sources \\citep{mereghetti95ApJ:AXPs}, however, \\citet{gavriil02Natur:AXPs} reported  in 2002 the discovery of SGR-like X-ray bursts from AXP 1E~1048.1$-$5937. Many other burst detections from AXPs followed (with the exception of giant flares), implying a possible evolutionary link between these two classes of magnetars \\citep[see][for reviews]{woods05ApJ:xte1810, mereghetti95ApJ:AXPs, perna11ApJ}.  Magnetars become burst active randomly with active periods lasting for relatively small time-intervals of weeks to months, between several years of quiescence. The total energy released during these episodes varies from source to source, and between active episodes of the same source \\citep[e.g., ][]{kouveliotou93Natur:1900}. Past spectral analyses of SGR short and intermediate bursts revealed that an optically thin thermal bremsstrahlung (OTTB) model fits the data above $\\sim15$~keV well, with temperatures in the range of 20-40~keV \\citep{aptekar01ApJS:SGR}. However, the OTTB model overestimates fluxes below 15 keV \\citep[e.g., ][]{fenimore94ApJ:1806, feroci04ApJ:1900}. Models consisting of a two black-body (2BB) or a power-law (PL) with a high-energy cutoff (Comptonized model; $f=A\\exp{[-E(2+\\lambda)/E_{\\rm peak}](E/E_{\\rm piv})^\\lambda}$, where $E_{\\rm pivot}$ is set to 30~keV) were shown to best fit the 1-150~keV broad-band spectra of SGR bursts \\citep[e.g., ][V12 hereafter]{feroci04ApJ:1900,olive04ApJ:1900, lin12ApJ:1550,vanderhorst12ApJ:1550}.  The origin and radiative mechanism of the persistent and burst emission from magnetars are not yet fully understood. In the magnetar framework, the decay of the magnetic field powers the quiescent emission through fracturing of the NS crust and sub-surface heating \\citep{thompson96ApJ:magnetar}. Bursts occur when a build-up of magnetic stresses on the NS crust from a twisted magnetic field causes the crust to crack, ejecting hot plasma into the magnetosphere \\citep{thompson95MNRAS:GF,thompson02ApJ:magnetars}. Alternatively, magnetar bursts could result from magnetic field reconnection in the magnetosphere \\citep[][see also \\citealt{gill10MNRAS:SGRGF}]{ lyutikov03MNRAS:recon}. Although the distinction between the two models is extremely challenging to measure, detailed statistical, temporal, and spectral studies of magnetar bursts have been effectively used for model comparisons \\citep[e.g.,][V12]{ gogus99apjl:1900, gogus00ApJ:1806, woods05ApJ:xte1810, israel08ApJ:1900, lin12ApJ:1550, lin11ApJ:0501, savchenko10AA:1547}.  \\src, which is the subject of this study, was suggested as a magnetar candidate in the supernova remnant G327.24$-$0.13 based on its X-ray spectral shape and varying flux \\citep{gelfand07ApJ:1547}. Subsequent radio observations revealed its magnetar nature with the discovery of radio pulsations at a period of 2.1~s, and $\\dot{P}=2.3\\times10^{-11}$~s$^{-1}$, implying a surface field strength of $2.2\\times10^{14}$~G \\citep{camilo07ApJ:1547}. A \\xmm\\ observation caught the source in outburst, which allowed the detection of X-ray pulsations at the same spin period \\citep{ halpern08ApJ:1547}. No bursts from \\src\\ were detected until 2008 when the source entered three active episodes: October 2008 \\citep{israel10MNRAS:1547, vonkienlin12ApJ:1550}, January 2009, and March 2009. The most prolific activity took place on 22 January 2009, when the source emitted hundreds of bursts in 24 hours. These latter bursts were detected by multiple high-energy instruments such as the \\fermi/GBM, \\swift/BAT, \\rxte/PCA, and {\\sl INTEGRAL}, and their temporal and spectral properties have been studied extensively \\citep[V12;][]{mereghetti09ApJ:1547,savchenko10AA:1547,kaneko10ApJ:1550,scholz11ApJ:1547}.  Here we report on the time-resolved spectroscopy of the brightest bursts from \\src\\ seen with \\fermi/GBM, which is ideal for such analyses owing to its high time-resolution and spectral capabilities. In particular, we concentrate on the spectral evolution of these bursts using the two spectral models established earlier as best for the source (V12), i.e., Comptonized model and the 2BB model. In Section~\\ref{datSel}, we describe the data selection and burst identification technique. We present our results on the time-resolved spectroscopy in Section~\\ref{timeResSpec}. In Section~\\ref{disc}, we first compare our results to the ones derived with time-integrated spectroscopy, and then to previous results on time-resolved spectroscopy of different sources. We interpret our results in the context of the current theoretical models for the origin of magnetar bursts, and the radiative processes that accompany them. Finally, we conclude with a few remarks in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper, we have shown the importance of high time-resolution in the study of the spectral evolution of SGR bursts. We were able to put to test the emission from a hot, optically thick, fireball of photon-pair plasma, predicted to form during SGR short bursts, according to a number of models for triggering magnetar flares (TD95, \\citealt{gill10MNRAS:SGRGF}, \\citealt{heyl05ApJ:magBurst}). We were able to follow the evolution of the emitting plasma in an unprecedented detail. Nonetheless, the very detailed observational picture presented in this paper motivates a more in depth modeling of the development and evolution of scattering/splitting cascades involving photon-pair fireballs in magnetar magnetospheres, taking into account the many physical and geometrical effects, to better understand the emission mechanism taking place during magnetar bursts."
    },
    "1402/1402.4821_arXiv.txt": {
        "abstract": "Astronomical phenomena are governed by processes on all spatial and temporal scales, ranging from days to the age of the Universe (13.8\\,Gyr) as well as from km size up to the size of the Universe. This enormous range in scales is contrived, but as long as there is a physical connection between the smallest and largest scales it is important to be able to resolve them all, and for the study of many astronomical phenomena this governance is present. Although covering all these scales is a challenge for numerical modelers, the most challenging aspect is the equally broad and complex range in physics, and the way in which these processes propagate through all scales.  In our recent effort to cover all scales and all relevant physical processes on these scales we have designed the Astrophysics Multipurpose Software Environment (AMUSE). AMUSE is a Python-based framework with production quality community codes and provides a specialized environment to connect this plethora of solvers to a homogeneous problem solving environment. ",
        "introduction": "The universe evolves over multiple scales in time and space.  This wide range in scales is generally hard to simulate on a standard 64-bit computer. If a causal connection between the scales exists multi-scale simulations are necessary.  In some cases, such a connection is made, for example in gravity, where the long-range and the absence of shielding causes the small scales to be connected with the largest scales. To the extreme, one can argue that the dynamical evolution of a star at the other side of the Milky Way Galaxy may affect life on the planet Earth. This works as follows: the star at the other side of the Milky Way may have been in the cluster in which the Solar system was born. Early perturbations in the solar system may than have originated from this nearby star, which in the last 4.6\\,Gyr has moved away to be at the other side of the Milky Way today. On the same account, a star that is in the neighborhood today may have been born in an entirely different galaxy, which merged several Gyr ago with the Milky Way Galaxy. In this way long ago, and far away events may have a direct imprint on the evolution of the solar system here and now, and with that on the evolution of the planet Earth.  In order to appreciate the complexity of the dynamical evolution of the Galaxy scales ranging in size from one to $10^{18}$\\,km ($10^5$\\,parsec), and from a single second to over $10^{17}$\\,seconds (4.6\\,Gyr) should be appreciated.  Incorporating such a wide range in scales in numerical simulations is extremely challenging, and luckily for many modern studies most scales can be omitted with the argument that the relative importance of mutual interactions between the smallest and the largest scales is small. In other cases this decoupling is not so evident, and such an argument should be verified explicitly.  On top of the different time and distance scales in our Universe, the physical processes happening on these scales have traditionally been split into different domains. Gravity underpins most of the processes in the Universe, and models have been developed to predict the typical evolution of gravitational bodies at different scales. From very specialized models that can predict the evolution of a small set of bodies \\((N<10^2)\\) \\cite{1999MNRAS.304..793C} \\cite{2008AJ....135.2398M}, and models that can accurately predict the dynamics of a modest number of bodies \\((10^2<N<10^4)\\) \\cite{2007NewA...12..357H} \\cite{1999PASP..111.1333A} to models that are less accurate but able to predict the overall evolution of many bodies \\((10^6<N<10^9)\\) \\cite{2005MNRAS.364.1105S} \\cite{2012JCoPh.231.2825B}. Star clusters, stars and planets are created from collapsing clouds of gas on different length and time scales. Like gravitational models, many hydrodynamical models for predicting the evolution of these gas clouds have been developed \\cite{2008ApJS..178..137S} \\cite{1989ApJS...70..419H} \\cite{1991A&A...252..718M}, each having unique properties that we would like to utilize. Similarly, to describe the evolution of a single star, a variation of stellar evolution models can be applied. From plasma magneto-hydrodynamical models for very short time and distance scales \\cite{2012JCoPh.231..718K}, and 1d spherical symmetrical models of a star for an accurate description of the processes inside a star over long time-scale \\cite{1971MNRAS.151..351E} \\cite{2008A&A...488.1007G} \\cite{2011ApJS..192....3P}, to analytically models that describe a few important variable of a star over very long time scales \\cite{2000MNRAS.315..543H} \\cite{1996A&A...309..179P}. Finally, radiative processes need to be modeled to predict the ionization levels and energy balance of the different gas clouds in the Universe. Specialized models exists for radiation on different bands like X-ray and UV, different time-scales and space scales \\cite{2010A&A...515A..79P} \\cite{2008MNRAS.386.1931A}.  In the past 4 years we have developed the Astrophysics Multipurpose Software Environment (AMUSE) to simulate parts of our Universe. Within AMUSE we want to include large collection of codes, to be agnostic to which implementation of a physical model we can to use for any simulation. Also, apart from the physical model, the actual algorithm often plays an important role in the accuracy and efficiency by which we can simulate that physical model. Finally, a large part of our predictive research is figuring out which physical processes play an important role in the system we are interested in and we do not want to be pinned to any specific model or algorithm.  We have designed AMUSE as a toolbox of different simulation codes with specialized support to couple these codes based on different scales and physics \\cite{2013CoPhC.183..456P} \\cite{2009NewA...14..369P} \\cite{2012ASPC..453..317P} \\cite{2008AN....329..885H}. In this paper we will discuss our experiences in the development of AMUSE, which is a framework for multi-scale and multi-domain physics. AMUSE is publicly available at {\\tt http://amusecode.org}. ",
        "conclusions": "It is challenging to computationally evolve an astrophysical system over multiple scales involving a variety of physical processes. We have tried to tackle this problem by creating a framework of astrophysical codes were all quantities have units, all data is represented in the same way and all codes provide standard components that support interaction between the codes. Since the coupling between the codes is implemented in the controlling layer of the framework, this approach depends on the possibility to split the physics of a system into separate parts either over the scales or over the type of physics involved or both. We have found that AMUSE is well suited for a range of problems in astrophysics where such a separation suggests itself from te nature of the problem. In other cases, where the physical processes are tightly intertwined it is best to import the physical solver in question as a single code. The exact boundary between these two cases remains an area of active research.  AMUSE is completely free and we distribute it as open source via the website {\\tt http://amusecode.org}."
    },
    "1402/1402.2293_arXiv.txt": {
        "abstract": "{Understanding solar surface magnetoconvection requires the study of the Fourier spectra of the velocity fields. Nowadays, observations are available that resolve very small spatial scales, well into the subgranular range, almost reaching the scales routinely resolved in numerical magnetoconvection simulations.  Comparison of numerical and observational data at present can provide an assessment of the validity of the observational proxies.} {Our aims are: (1) to obtain Fourier spectra for the photospheric velocity fields using the spectropolarimetric observations with the highest spatial resolution so far ($\\sim 120$ km), thus reaching for the first time spatial scales well into the subgranular range; (2) to calculate corresponding Fourier spectra from realistic 3D numerical simulations of magnetoconvection and carry out a {{\\it proper comparison}} with their observational counterparts considering the residual instrumental degradation in the observational data; and (3) to test the observational proxies on the basis of the numerical data alone, by comparing the actual velocity field in the simulations with synthetic observations obtained from the numerical boxes.} {(a) For the observations, data from the SUNRISE/IMaX spectropolarimeter are used. (b) For the simulations, we use four series of runs obtained with the STAGGER code for different average signed vertical magnetic field values ($0$, $50$, $100$, and $200$ G). Spectral line profiles are synthesized from the numerical boxes for the same line observed by IMaX (Fe I 5250.2 \\AA) and degraded to match the performance of the IMaX instrument. Proxies for the velocity field are obtained via Dopplergrams (vertical component) and LCT (local correlation tracking, for the horizontal component). Fourier power spectra are calculated and a comparison between the synthetic and observational data sets carried out. (c) For the internal comparison of the numerical data, velocity values on constant optical depth surfaces are used instead of on horizontal planes.} {A very good match between observational and simulated Fourier power spectra is obtained for the vertical velocity data for scales between $200$ km and $6$ Mm. Instead, a clear vertical shift is obtained when the synthetic observations are not degraded to emulate the degradation in the IMaX data. The match for the horizontal velocity data is much less impressive because of the inaccuracies of the LCT procedure. \\ Concerning the internal comparison of the direct velocity values of the numerical boxes with those from the synthetic observations, a high correlation ($0.96$) is obtained for the vertical component when using the velocity values on the $\\log \\tau_{500} = -1$ surface in the box. The corresponding Fourier spectra are near each other. A lower maximum correlation ($0.5$) is reached (at $\\log \\tau_{500} = 0$) for the horizontal velocities as a result of the coarseness of the LCT procedure. Correspondingly, the Fourier spectra for the LCT-determined velocities is well below that from the actual velocity components.} {As measured by the Fourier spectra, realistic numerical simulations of surface magnetoconvection provide a very good match to the observational proxies for the photospheric velocity fields at least on scales from several Mm down to around $200$ km. Taking into account the spatial and spectral instrumental blurring is essential for the comparison between simulations and observations. Dopplergrams are an excellent proxy for the vertical velocities on constant-$\\tau$ isosurfaces, while LCT is a much less reliable method of determining the horizontal velocities.} ",
        "introduction": "\\label{sec:introduction}  The convective flows in the outermost layers of the solar interior and photosphere play an essential role in the dynamics and magnetism of the solar envelope and atmosphere. Their kinetic energy is distributed over a very wide range of space- and timescales \\citep[see, e.g.,][]{nordlund2009}, as results from (a) the strong stratification in the topmost tens of Megameters below the surface, leading to the appearance of  convection cells from granular to supergranular size; and (b) the very high Reynolds numbers of those flows, which lead to turbulence and cascading of energy from the granular size downward to small scales before reaching the dissipation range.  Solar convection provides a prime example of the need for interaction between observations and numerical simulations.  The most recent advances in observational facilities now permit the observation of structures, flows, and fields with very high spatial resolution,  down to $\\approx 100$ km. On the other hand, the numerical experiments and simulations of the past ten years, using the power of recent supercomputing installations, have reached high realism concerning the physical processes studied in the solar interior and photosphere (including material properties and detailed radiative transfer) with numerical grid spacing down to the order of $10$ km \\citep[e.g.,][]{2006ApJ...642.1246S, 2007A&A...465L..43V, 2008ApJ...684L..51J}. In addition, for the large scales, realistic simulations are now being produced on boxes with sides spanning several tens of Megameters in the horizontal directions, hence including a few supergranular cells in them \\citep[e.g.,][]{2011SoPh..268..271S}. The simulations of convection with different numerical codes seem to agree reasonably well \\citep{Beeck_etal_2012}.  A crucial tool to understanding the nature of the photospheric flows is their Fourier power spectrum.  Application of this tool has been made along the years to both observational and numerical simulation data, mostly separately, but sometimes also comparing results from the two sources \\citep[e.g.,][]{Stein_Nordlund1998, 2006ASPC..354..109G, Georgobiani2007, Kitiashvili_etal_2012}. For the observations, proxies have to be used for the velocity components, like Doppler shifts of spectral lines for the vertical velocity \\citep[e.g.,][]{Hathaway_etal_2000, Stein_etal_2006, Georgobiani2007, rieutord2010, Kitiashvili_etal_2012, Katsukawa2012} or feature displacement on series of 2D continuum or Doppler maps \\citep[e.g.,][]{rieutord2001, rieutord2010, Stein_etal_2006, 2006ASPC..354..109G, Georgobiani2007, 2012A&A...540A..88R, Goode_etal_2010, Abramenko2012, Kitiashvili_etal_2012}. For the simulations, the velocity values on horizontal cuts in the numerical box have been Fourier-analyzed \\citep[][]{Stein_Nordlund1998, rieutord2001, Stein_etal_2006, 2006ASPC..354..109G, Georgobiani2007, Kitiashvili_etal_2012}.  In the current paper we would like to push the limits of the comparative analysis between simulations and observations in  three directions:  \\begin{enumerate}[(i)]  \\item \\label{enum:i} by using Fourier spectra from observational proxies for the photospheric velocity field (Doppler, Local Correlation Tracking) using as a basis recent observations with very high spatial resolution;  \\item \\label{enum:ii} by comparing them with Fourier spectra from numerical simulations but taking great care that the comparison is between homogeneous datasets, i.e., by first obtaining {\\it synthetic observations} from the simulations and then modifying them to reproduce the instrumental degradation suffered by the real observations; and  \\item \\label{enum:iii} by checking the closeness of the synthetic observational proxies to the actual velocity vector components in the numerical box when using surfaces of constant optical depth ($\\tau_{500}$=const) that correspond to the formation region of the spectral lines.  \\end{enumerate}   For the first aspect, observations with spatial resolution as high as one-tenth of the average granular size are available at present.  Using Hinode/SOT data (spatial resolution: $200$ km), \\cite{rieutord2010} have explored the power spectra of the vertical velocity down to scales of $400$ km (and of the horizontal velocity in a range from $2.5$ Mm to $76.3$ Mm), but no comparison to numerical simulations was provided. \\citet{Goode_etal_2010} obtained continuum images providing information down to very small scales of about $90$ km using speckle-reconstruction techniques with the NST/BBSO telescope. From them, they could obtain horizontal-velocity maps using Local Correlation Tracking (LCT) techniques. In the present paper we take advantage of the availability of time series of spectropolarimetric observations with unprecedented spatial resolution (near $120$ km) obtained using the IMaX instrument aboard the SUNRISE balloon mission \\citep{martinezpillet_etal_2011, solanki_etal_2010}. These data allowed the determination of the line-of-sight velocity, continuum maps (used to compute the horizontal velocity by applying LCT) and of the full magnetic field vector with a field of view of $32 \\times 32$ Mm$^2$ and for periods of half an hour with a $33.3$-sec cadence.      For the second aspect, comparisons so far have been done almost exclusively by using observational proxies, on the one hand, and 2D maps on horizontal planes in the numerical box, on the other. Yet, as done in the present paper, using spectral synthesis it is possible to obtain synthetic observations from the simulations that can then be subjected to the same sort of quality degradation that the solar spectrum suffers when going through the observational equipment: for that, one only needs to carry out a convolution of the synthetic spectra with a Point Spread Function (PSF) in the spatial domain and with the Monochromatic Spectral Response Function (MSRF) in the spectral domain adequate to the specific equipment used to obtain the observations. In this way, the comparison can be carried out using the same kind of proxies on either side.    Finally, for the third aspect, the numerical datasets and techniques used in this paper allow us to approach the question of the closeness of the observational proxies (Dopplergrams and horizontal-velocity maps from feature tracking) to the actual velocity fields.  So, it is of interest to check if the observational proxies are near the actual velocity field when both are obtained from the numerical results.  The actual velocity field data have to be taken from a representative surface in the box, for which we choose an isosurface of the optical depth at $500$ nm, $\\tau_{500}$,  which seems more appropriate than a horizontal plane in the numerical box.  This is because, following basic spectral line formation theory \\citep[e.g.,][]{gray2005}, the lines used for the Doppler velocity determination here are expected to be formed roughly in the same $\\tau_{500}$--range for all vertical columns in the numerical box or in the observations (see further explanations in Sect.~\\ref{sec:velocities_at_grid_points}).  The layout of the paper is as follows: Sect.~\\ref{sec:observations} and \\ref{sec:simulations} describe the observational and numerical datasets, respectively, used in the paper. Section~\\ref{sec:IMaXconditions} details the procedure whereby we gain synthetic observations from the numerical boxes and the instrumental degradation that we have to apply to make these synthetic observations coherent with the observational data. The core of the paper are Sect.~\\ref{sec:vertical} through \\ref{sec:velocities_at_grid_points}. Sect.~\\ref{sec:vertical} and \\ref{sec:horizontal} discuss the Fourier power spectra obtained for the vertical and horizontal velocity distributions, respectively. Section~\\ref{sec:velocities_at_grid_points} compares the observational proxies with the numerical data taken on $\\tau_{500}$ isosurfaces. Sections~\\ref{sec:summary} and \\ref{sec:conclusions} provide a discussion and conclusions. ",
        "conclusions": "\\label{sec:conclusions}    As measured using Fourier power spectra, realistic numerical simulations of surface magnetoconvection provide an accurate picture of the photospheric vertical velocity fields in the range from a few $100$ km to several Mm. This conclusion is based on a comparison of Dopplergrams obtained from real observations, on the one hand, and from synthetic observations calculated from the numerical data, on the other. Using as observational proxy for the horizontal velocities the Local Correlation Tracking method, the match is worse due to the limitations of that technique.  Taking detailed care of spatial and spectral blurring caused by the observational equipment is necessary for a proper comparison between numerical and observational data: the instrumental degradation can seriously affect the Fourier power spectra of the LOS velocities (compare Figure~\\ref{fig2} with Figure~\\ref{fig4}), leading to a much worse match of synthetic and real observations. Apodization of the observational data prior to Fourier processing is advisable but less crucial.  We have compared the vertical velocity proxy (Dopplergrams) in the synthetic observations with the actual velocity field on the individual $\\tau_{500} =$ const isosurfaces in the numerical box. The maximum correlation ($0.96$) is found for the $\\log\\tau_{500}=-1$ surface, which is $\\sim 200$ km above the continuum-forming layer ($\\log\\tau_{500}=0$). The corresponding Fourier spectra are very close to each other. The typical granulation velocity in those heights would be some $0.8-0.9$ km/s (Figure~\\ref{fig8}).  A similar comparison for the horizontal velocities leads to a smaller maximum correlation ($0.5$), which is reached at the $\\tau_{500}=1$ isosurface and to the spectrum for the actual velocity field on that isosurface being clearly above the spectrum for the LCT-determined velocity proxy (Figure~\\ref{fig9}). Both problems can be ascribed to the limitations of the LCT technique particularly when dealing with small-scale features.  Comparing cases with different magnetic flux, we tentatively conclude that a higher average vertical flux leads to less kinetic energy and to a shift in relative weight of the different scales toward the small-$\\lambda$ range."
    },
    "1402/1402.2770_arXiv.txt": {
        "abstract": "Chemical models predict that the deuterated fraction (the column density ratio between a molecule containing D and its counterpart containing H) of \\H , \\Dfrac (\\H ), high in massive pre-protostellar cores, is expected to rapidly drop of an order of magnitude after the protostar birth, while that of HNC, \\Dfrac (HNC), remains constant for much longer. We tested these predictions by deriving \\Dfrac (HNC) in 22 high-mass star forming cores divided in three different evolutionary stages, from high-mass starless core candidates (HMSCs, 8) to high-mass protostellar objects (HMPOs, 7) to Ultracompact HII regions (UCHIIs, 7). For all of them, \\Dfrac (\\H ) was already determined through IRAM-30m Telescope observations, which confirmed the theoretical rapid decrease of \\Dfrac(\\H ) after protostar birth (Fontani et al. 2011). Therefore our comparative study is not affected by biases introduced by the source selection. We have found average \\Dfrac (HNC) of 0.012, 0.009 and 0.008 in HMSCs, HMPOs and UCHIIs, respectively, with no statistically significant differences among the three evolutionary groups. These findings confirm the predictions of the chemical models, and indicate that large values of \\Dfrac (\\H ) are more suitable than large values of \\Dfrac (HNC) to identify cores on the verge of forming high-mass stars, likewise what found in the low-mass regime. ",
        "introduction": "\\label{intro}  The process of deuterium enrichment in molecules from HD, the main reservoir of deuterium in molecular clouds, is initiated by three exothermic ion-molecule reactions (e.g.~Millar et al.~1989): \\begin{equation} {\\rm H_3^+ + HD \\rightarrow H_2D^+ + H_2}\\;, \\label{eqa} \\end{equation} \\begin{equation} {\\rm CH_3^+ + HD \\rightarrow CH_2D^+ + H_2}\\;, \\label{eqb} \\end{equation} \\begin{equation} {\\rm C_2H_2^+ + HD \\rightarrow C_2HD^+ + H_2} \\;. \\label{eqc} \\end{equation} Since the backward reactions are endothermic by 232, 390 and 550 K, respectively, in cold environments (e.g.~$T_{\\rm kin} \\leq 20$ K for reaction (1)) they proceed very slowly, favouring the formation of deuterated ions. Moreover, the freeze-out of CO and other neutrals, particularly relevant in high-density gas ($n_{\\rm H_2}\\geq 10^4$ \\cmc ), further boosts the deuteration process (e.g.~Bacmann et al.~\\citeyear{bacmann03}, Crapsi et al.~\\citeyear{crapsi05}, Gerin et al.~\\citeyear{gerin06}). Therefore, in dense and cold cores the deuterated fraction, \\Dfrac , defined as the abundance ratio between a deuterated molecule and its hydrogenated counterpart, is expected to be much higher than the average [D/H] interstellar abundance (of the order of $10^{-5}$, Oliveira et al.~\\citeyear{oliveira03}, Linsky et al. ~\\citeyear{linsky06}). Because of the changements in physical and chemical properties of a star forming core, its \\Dfrac\\ is expected to change with the evolution too. Specifically, \\Dfrac\\ is predicted to increase when a pre--protostellar core evolves towards the onset of gravitational collapse, as the core density profile becomes more and more centrally peaked (due to the temperature decrease at core centre, e.g.~Crapsi et al.~\\citeyear{crapsi07}), and then it drops when the young stellar object formed at the core centre begins to heat its surroundings (see e.g. Caselli et al.~\\citeyear{caselli02}).  While this net drop in \\Dfrac\\ before and after the protostellar birth in \\Dfrac (\\H ) is clearly observed in both low-mass (Crapsi et al.~\\citeyear{crapsi05}, Emprechtinger et al.~\\citeyear{emprechtinger09}) and high-mass (Fontani et al.~\\citeyear{fontani11}, Chen et al.~\\citeyear{chen11}) star-forming cores, other species show deviations from this general scenario. For example, DNC is produced in the gas from the same route reaction as \\D , namely reaction (\\ref{eqa}), so that \\Dfrac (HNC) and \\Dfrac (\\H ) are expected to vary similarly with temperature (Turner~\\citeyear{turner}). However, Sakai et al.~(\\citeyear{sakai12}) have measured \\Dfrac (HNC) in a sample of 18 massive cores including both infrared-dark starless cores and cores harbouring high-mass protostellar objects, and found that \\Dfrac (HNC) in the starless cores is only marginally higher than that measured in the protostellar cores. This 'anomaly' could be explained by the fact that the destruction processes of \\D\\ are much faster than those of DNC: being an ion, \\D\\ can recombine quickly (few years) with CO and/or electrons, while the neutral DNC has to be destroyed by ions (such as \\HCO\\ and/or H$_3^+$) through much slower ($10^4 - 10^5$ yrs) chemical reactions (Sakai et al.~\\citeyear{sakai12}).  The chemical models of Sakai et al.~(\\citeyear{sakai12}), are able to partially reproduce the observational results obtained by Fontani et al.~(\\citeyear{fontani11}) and Sakai et al.~(\\citeyear{sakai12}). We have performed chemical calculations similar to those in Sakai et al.~(\\citeyear{sakai12}), and obtained the consistent results (see Section 4.2 for the details of our chemical model): the \\D /\\H\\ ratio approaches $\\sim 0.1$ during the cold pre--protostellar phase and drops quickly to $\\sim 0.01$ after the protostellar birth because very sensitive to a temperature growth (Fig.~1, panel (a)), while the DNC/HNC ratio remains relatively high (above $\\sim 0.01$) even after a rapid temperature rise, and decreases in timescales of several 10$^4$ yrs. On the other hand, the \\D /\\H\\ drops much more quickly, in less than 100 yrs (Fig.~1, panel (b)). However, the different criteria adopted by Fontani et al.~(\\citeyear{fontani11}) and Sakai et al.~(\\citeyear{sakai12}) to select the targets does not allow for a consistent observational comparison between the two deuterated fractions, as well as between models and data.  \\begin{figure*} \\begin{minipage}{160mm} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{fig1.eps}} \\caption[] {\\label{models}{{\\bf (a)}: chemical model calculation of the time dependence of DNC/HNC (thick lines) and N$_2$D$^+$/N$_2$H$^+$ (thin lines).  In addition to the constant temperature cases (10 K = solid black, 30 K = solid grey), the cases in which the temperature suddenly rises from 10 to 30 K at a given age, 1$\\times$10$^5$ yr (dashed) or 3$\\times$10$^5$ yr (dotted) are shown. See Section~\\ref{model} for the detail of the chemical model. \\newline {\\bf (b)}: same as panel {\\bf (a)}, but the time starts from when the temperature increases from 10 K to 30 K. }} \\end{center} \\end{minipage} \\end{figure*}  In this paper we report observations performed with the Nobeyama-45m Telescope in the DNC and HN$^{13}$C (1--0) rotational transitions towards 22 high-mass cores harbouring different stages of the high-mass star formation process, in which \\Dfrac (\\H ) was already measured through observations of the IRAM-30 Telescope (Fontani et al.~\\citeyear{fontani11}). In this way, our study is not affected by observational biases possibly introduced by the source selection. The main aim of the work is to test in the same sample of objects whether the \\D /\\H\\ and DNC/HNC ratios trace differently the thermal history of high-mass cores despite the similar chemical origin, as predicted by Sakai et al.~(\\citeyear{sakai12}). The sample, selected as explained in Fontani et al.~(\\citeyear{fontani11}), is divided in 8 high-mass starless cores (HMSCs), 7 high-mass protostellar objects (HMPOs) and 7 ultracompact HII regions (UC HIIs), so that all main evolutionary groups of the high-mass star formation process are almost equally represented. We stress that all HMSCs, except I22134--B, have been previously classified as 'quiescent' by Fontani et al.~(\\citeyear{fontani11}) to distinguish them from 'perturbed' cores, in which external phenomena (passage of outflows, shocks, nearby infrared objects) can have affected significantly the physical-chemical properties of the gas, as discussed in Fontani et al.~(\\citeyear{fontani11}).  In Sect.~\\ref{obs} we give an overview of the technical details of the observations; Sect.~\\ref{res} presents the main observational results, which are discussed in Sect.~\\ref{dis}, including a detailed comparison with chemical models. A summary of the main findings of the paper are given in Sect.~\\ref{con}. ",
        "conclusions": "\\label{con}  We have observed the DNC and \\HNC (1--0) rotational transitions towards 22 massive star-forming cores in different evolutionary stages, towards which \\Dfrac (\\H) was already measured by Fontani et al.~(\\citeyear{fontani11}). The aim of the work was to compare \\Dfrac (HNC) to \\Dfrac (\\H ) in the same sample of sources and with similar telescope beams, so that the comparison should not suffer from possible inconsistencies due to different sample selection criteria. The main observational result of this work confirms the predictions of the models of Sakai et al.~(\\citeyear{sakai12}), namely that \\Dfrac (HNC) is less sensitive than \\Dfrac (\\H ) to a sudden temperature rise, and hence it should keep more than \\Dfrac (\\H ) of the thermal history of the cores, despite the chemical processes leading to the deuteration of the two species are similar. Therefore, our work clearly indicates that \\Dfrac (\\H ) is more suitable than \\Dfrac (HNC) to identify high-mass starless cores. Based on the predictions of our chemical models, the starless cores studied in this work having \\Dfrac (\\H ) around 0.2 - 0.3 are very good candidate massive 'pre--stellar' cores, because only relatively 'evolved' starless cores can be associated with such high values of \\Dfrac (\\H ). Several results require follow-up high-angular resolution observations to map the emitting region of DNC and \\D , as well as DNC and \\HNC , and test if these are slightly different as the present low-angular resolution data seem to suggest. Observations of higher excitation \\HNC\\ and DNC lines (3--2 or 4--3) may be also important to constrain better the excitation conditions."
    },
    "1402/1402.6259_arXiv.txt": {
        "abstract": "We investigate the production of antideuterons from decaying dark matter, using gravitinos in supersymmetric models with trilinear R-parity violating (RPV) operators as an example. The model used for antideuteron formation is shown to induce large uncertainties in the predicted flux, comparable to uncertainties from cosmic-ray propagation models. We improve on the formation model by tuning hadronization and coalescence parameters in Monte Carlo simulations to better reproduce the hadron spectra relevant for antideuteron production. In light of current bounds on fluxes and future prospects from the AMS-02 and GAPS experiments we set limits on RPV couplings as a function of the gravitino mass. ",
        "introduction": "The last few years have seen a lot of interest in decaying dark matter scenarios, in which dark matter (DM) candidates are very long lived particles. Having decaying rather than annihilating DM candidates leads to important consequences for indirect detection signals, as the flux of decay products scale with the density of dark matter, whereas the flux of annihilation products scales with the DM density squared. Such long lived particles would be natural if the only coupling to Standard Model particles is very weak. Gravitational interaction is one such alternative.  In supersymmetric models, negative results from LHC searches for the standard missing energy signature of dark matter suggest looking at alternative models where R-parity is violated. With the gravitino as the lightest supersymmetric particle, the strength of the gravitational coupling means that gravitinos can naturally have lifetimes exceeding the age of the Universe. Such gravitino DM scenarios have been considered for some time now, beginning with~\\cite{Takayama:2000uz,Buchmuller:2007ui,Lola:2007rw}, with specific indirect detection signatures of gamma-rays, positrons, antiprotons and neutrinos.  With a well motivated example of decaying dark matter, we will here investigate the production of antideuterons in gravitino decays for models with trilinear R-parity violating (RPV) operators. Use of the antideuteron channel for indirect detection of DM was first suggested in Ref.~\\cite{Donato:1999gy}, and antideuterons have already been considered for generic decaying dark matter~\\cite{Ibarra:2009tn}, and also specifically for gravitino DM~\\cite{Grefe:2011dp} in models with bilinear RPV operators. We note several important differences between the present paper and previous works: The generic DM study~\\cite{Ibarra:2009tn} includes only two-body decays with no baryon or lepton number violation, and therefore cannot describe the three-body decays of gravitinos into quarks and leptons due to  trilinear RPV operators. The work on RPV gravitino decays~\\cite{Grefe:2011dp} does take three-body RPV decays into account, but only for models with bilinear RPV operators. We shall see that the baryon number violating trilinear couplings produce significantly more antideuterons than any other RPV decays previously considered. Moreover, we will use the resulting predicted antideuteron flux at Earth to set limits on the trilinear RPV couplings, and unlike~\\cite{Ibarra:2009tn}, we do not make the simplifying assumption of isotropic and uncorrelated distributions of antiprotons and antineutrons in the decays.  However, for readers interested in antideuteron limits for generic dark matter models, the most interesting novelty in this paper is probably the inclusion and treatment of uncertainties in the formation model for antideuterons, in particular hadronization effects. We investigate the effect of phenomenological hadronization parameters in the {\\tt Herwig++} Monte Carlo event generator~\\cite{Bahr:2008pv,Arnold:2012fq} and the $p_0$ parameter of the coalescence model, and re-tune these to better describe the particular hadron spectra relevant for antideuteron production. We present a set of recommended parameters for further studies that result in an improved description of antideuteron formation with smaller and quantifiable uncertainties.  The trilinear RPV terms involved in the gravitino decays under study have the form \\beq W\\sim\\lambda_{ijk} L_iL_j\\bar{E}_k + \\lambda'_{ijk}L_iQ_j\\bar{D}_k+\\lambda''_{ijk}\\bar{U}_i\\bar{D}_j\\bar{D}_k, \\label{eq:WRPV} \\eeq in the superpotential. Here $\\lambda_{ijk}$, $\\lambda'_{ijk}$ and $\\lambda''_{ijk}$ are RPV couplings, $Q$ and $L$ are left-handed quark and lepton superfield doublets respectively, while $\\bar{E}$, $\\bar{U}$ and $\\bar{D}$ are the corresponding left-handed superfield singlets for leptons and up and down type quarks. The first two terms violate lepton number, while the third violates baryon number conservation. From gauge symmetry, $i\\neq j$ for $\\lambda_{ijk}$ and $j\\neq k$ for $ \\lambda''_{ijk}$. For the present purpose of antideuteron production we are interested in the couplings that give antiquarks in gravitino decays: $\\lambda'$ and $\\lambda''$.  While these operators, if allowed, can cause problems, most seriously with proton decay, completely removing them by introducing R-parity~\\cite{Fayet:1977yc} by hand seems excessive and {\\it ad hoc}. In particular it is known that the proton can be protected from decay by up to dimension--6 operators by introducing so-called trialities instead, disallowing either lepton or baryon number violation terms in Eq.~(\\ref{eq:WRPV})~\\cite{Ibanez:1991hv,Ibanez:1991pr,Lola:1993ip}.  Limits on trilinear couplings have been set on the basis of indirect detection experiments in the antiproton, positron and gamma-ray channels~\\cite{Lola:2007rw,Bomark:2009zm}. These are generically much stronger than other (collider) limits on RPV couplings, resting only on the assumption that the gravitinos constitute a significant fraction of the total DM. The limits' dependence on sparticle masses is dominated by the gravitino mass; the other sparticles are less important, unless there is a dramatic difference in scales compared to the gravitino mass~\\cite{Lola:2007rw}.\\footnote{For the purposes of this paper we set all other sparticle masses to 1 TeV, just outside the current reach of the LHC.}  As noted in~\\cite{Ibarra:2012cc}, DM decay and annihilation involving final state quark--antiquark pairs are in general already severely constrained by antiproton searches. However, we will see that the trilinear operators we investigate here have stronger bounds from antideuterons due to the possibility of a triple antiquark final state. Antideuterons from DM decay or annihilation are also expected to have a much flatter spectrum at low energies than those from the main background, cosmic-ray--gas interactions. In conjunction with the extremely low level of background expected, this  leads to even stronger limits on the trilinear couplings.  The present upper bound on the antideuteron flux that we use to set limits comes from the BESS balloon experiment~\\cite{Fuke:2005it}, while the AMS-02 experiment running onboard the ISS \\cite{AMSICRC30} and the planned GAPS balloon experiment~\\cite{Koglin:2008zz} are expected to deliver stronger bounds in the near future. We give our predictions for future limits given null results from these searches.  The starting-point of the calculation of the expected cosmic-ray flux of antideuterons from gravitino decays is the decay width for gravitinos in the processes $\\tilde G\\to \\nu_i d_j \\bar d_k,\\ell_i^- u_j \\bar d_k$ and $\\tilde G\\to u_i d_j d_k$ found in \\cite{Moreau:2001sr}. The parton level processes are showered and hadronized using the {\\tt Herwig++ 2.6.0} Monte Carlo event generator~\\cite{Bahr:2008pv,Arnold:2012fq}. The production of antideuterons from the resulting antiprotons and antineutrons is described by the so-called coalescence model. The complete formation model, from parton level to antideuterons, will be discussed in more detail in Sec.~\\ref{sec:Coalescence}, including the uncertainties from the hadronization model used. The derived antideuteron spectrum is then propagated through the Galaxy, and the antideuteron flux at Earth is found in Sec.~\\ref{sec:prop}. In Sec. \\ref{sec:limits} we then use this flux to set limits on the relevant trilinear RPV couplings, and make predictions for what limits can be set with future experiments. ",
        "conclusions": "In this paper we have studied the influence of coalescence and hadronization parameters on the yield of antideuterons in dark matter decay/annihilation. By fitting to relevant collider data we provide a recommended set of parameter values for the {\\tt Herwig++} event generator in Table~\\ref{tab:fitresults}, with corresponding uncertainties that can be used for error propagation.  With the tuned formation model we have set bounds on trilinear RPV couplings in scenarios with gravitino dark matter using the results of antideuteron searches with the BESS experiment. These are of the same order of magnitude as limits from antiproton searches with PAMELA. We find large differences between different types of operators, the antideuteron searches are far more constraining---an order of magnitude---for the baryon number violating operators, and are more constraining for operators with light quarks. We have further investigated the potential of AMS-02 and GAPS to improve on these limits, and found that future results can strengthen current limits by a factor of $\\sim 4$ for the baryon number violating operators in the whole range of gravitino masses considered (10--800 GeV). This will significantly improve on the current best limits for trilinear RPV couplings that come from antiproton searches."
    },
    "1402/1402.4776_arXiv.txt": {
        "abstract": "\\noindent Axions and similar very weakly interacting particles are increasingly compelling candidates for the cold dark matter of the universe. Having very low mass and being produced non-thermally in the early Universe, axions feature extremely high occupation numbers. It has been suggested that this leads to the formation of a Bose-Einstein condensate with potentially significant impact on observation and direct detection experiments. In this note we aim to clarify that if Bose-Einstein condensation occurs for light and very weakly interacting dark matter particles, it does not happen in thermal equilibrium but is described by a far-from-equilibrium state. In particular we point out that the dynamics is characterized by two very different timescales, such that condensation occurs on a much shorter timescale than full thermalization. ",
        "introduction": "Although dark matter makes up approximately one quarter of the total energy content of the Universe we still know very little about it. It is most likely not Standard Model matter and it interacts with Standard Model particles as well as with itself only very weakly. Moreover, it behaves much like a gas of non-relativistic particles, i.e. particles with very little kinetic energy, hence cold dark matter.  Therefore, it is not surprising that the two perhaps best known (and best motivated) candidates for dark matter particles, axions and weakly interacting massive particles (WIMPs), have very different properties. Axions are very light, sub-eV bosons whereas (SUSY) WIMPs are heavy, multi-GeV fermions. Axions are produced non-thermally from coherent oscillations of the axion field~\\cite{Preskill:1982cy,Abbott:1982af,Dine:1982ah} far from thermal equilibrium, whereas WIMPs are produced in thermal collisions in or close to thermal equilibrium. Nevertheless, both are thought to behave in many ways like a gas of cold, non-relativistic particles.  However, recently it has been suggested~\\cite{Sikivie:2009qn} that there is indeed a property that sets axions apart from WIMPs: Axions are bosons and if they are dark matter their low momentum states feature very high occupation numbers. This suggests the possibility that they form a Bose-Einstein condensate. The resulting collective behavior may lead to distinct observational signatures~\\cite{Sikivie:2012gi,Banik:2013rxa} as well as having important effects for experiments for the direct detection of axion dark matter~\\cite{Duffy:2006aa,Irastorza:2011gs,Jaeckel:2013sqa}. Usually Bose-Einstein condensation is treated in thermal equilibrium, and indeed the arguments of Refs.~\\cite{Sikivie:2009qn,Erken:2011vv,Erken:2012dz,Saikawa:2012uk} are based on estimates for thermalization rates\\footnote{Indeed the authors of the more recent papers~\\cite{Davidson:2013aba,Noumi:2013zga} more carefully refer to dissipation rates and self interaction rates.}. However, the extremely weak self-interactions of axions suggest that it takes very long times to reach full thermal equilibrium and establish a proper thermal distribution.  In this note we want to clarify this apparent conflict and argue that, while full thermal equilibrium\\footnote{We are considering only axions. When talking about full thermal equilibrium we mean thermal equilibrium of the axions by themselves. Our arguments do not concern the question of a possible equilibration with other species such as photons.} is probably not reached, a non-thermal Bose-Einstein condensate may nevertheless be formed. Based on results of non-equilibrium field theory~\\cite{Berges:2012us,Micha:2004bv,Nowak:2013juc,BS}, we elucidate the non-equilibrium nature of such a condensate: \\begin{itemize} \\item{} We clearly distinguish between the very different timescales for condensation $t_{\\rm cond}$ and that for thermalization $t_{\\rm therm}$ with $\\tcond\\ll\\ttherm$. \\item{} We estimate the time needed to reach a fixed value of the condensate fraction $N_0/N_{\\rm total}$ depending on the spatial volume. \\item{} The non-thermal condensate fraction can differ from that in thermal equilibrium. More importantly we note that the condensate is in a strongly correlated state which significantly deviates from a weakly interacting gas approximated by free field theory. \\end{itemize}  \\bigskip  Qualitatively, similar behavior may happen for other very light bosonic dark matter candidates such as axion-like particles (ALPs) or hidden photons~\\cite{Arias:2012ly} (see~\\cite{Jaeckel:2010ni} for a review of such particles).  For the purposes of this work the main feature of axions is that they are a very light bosonic particles with extremely weak self-interactions (we mostly neglect the similarly weak interactions with other matter particles). Therefore, our treatment should apply directly to axions and axion-like particles but similar mechanisms may apply for more general very light bosonic dark matter particles as well. For simplicity we will nevertheless refer to our particles as ALPs.  In the following we will only consider ALPs interacting via their self-interactions. As already found in~\\cite{Sikivie:2009qn} these interactions are probably not sufficient to achieve Bose-Einstein condensation of ALPs in the early Universe. The phenomenologically more relevant interactions are probably those via gravitons~\\cite{Sikivie:2009qn,Erken:2011vv,Erken:2012dz,Saikawa:2012uk,Davidson:2013aba,Noumi:2013zga}. Therefore our toy model does not allow to address the quantitative question whether the gravitational interactions of ALPs are fast enough to achieve condensation (both positive~\\cite{Sikivie:2009qn,Erken:2011vv,Erken:2012dz,Saikawa:2012uk,Noumi:2013zga} and negative\\cite{Davidson:2013aba} evidence exists). Importantly however, we expect that our general arguments hold and the far from equilibrium nature with the features described above will persist if such a Bose-Einstein condensate is formed via gravitational interactions. ",
        "conclusions": "We have analyzed the question of Bose-Einstein condensation for axion-like particles (ALPs) via their self-couplings. Our main findings are \\begin{itemize} \\item{} Thermalization, i.e. establishing a true thermal distribution with typical occupation numbers of order one for $p^{2}/(2m) \\sim T$, is  slow for weak coupling, even if we have extremely high initial occupation numbers. \\item{} The formation time for a condensate $\\tcond$ can be much smaller than the time required for thermalization $\\ttherm$. The corresponding timescale is similar (albeit a bit longer) than that obtained from an estimate of the Bose-enhanced interaction rates. \\item{} Condensate formation takes place in a non-equilibrium state and the condensate formed in this state is parametrically different from that obtained for a weakly interacting gas in thermal equilibrium. \\end{itemize}  In this note we have only considered scalar self-interactions. As already pointed out in~\\cite{Sikivie:2009qn} these interactions are not sufficient when it comes to a condensate formation in the expanding early Universe. Yet, importantly it can serve as a simple model for the basic dynamics of the formation of a Bose-Einstein condensate of ALP dark matter by Bose-enhanced interactions. Our positive finding that condensate formation takes indeed place on the corresponding short timescale now motivates studying the phenomenologically more relevant interactions via gravity."
    },
    "1402/1402.4295_arXiv.txt": {
        "abstract": "COupled SLope and scIntillation Detection And Ranging (CO-SLIDAR) is a recent profiling method of the vertical distribution of atmospheric turbulence strength ($C_n^2$ profile). It takes advantage of correlations of slopes and of scintillation, both measured with a Shack-Hartmann wavefront sensor on a binary star. In this paper, we present the improved CO-SLIDAR reconstruction method of the $C_n^2$ profile and the first on-sky results of the CO-SLIDAR profiler. We examine CO-SLIDAR latest performance in simulation, taking into account the detection noise bias and estimating error bars along with the turbulence profile. The estimated $C_n^2$ profiles demonstrate the accuracy of the CO-SLIDAR method, showing sensitivity to both low and high altitude turbulent layers. CO-SLIDAR is tested on-sky for the first time, on the $1.5$~m MeO (M\u00e9trologie Optique) telescope at Observatoire de la C\u00f4te d'Azur (France). The reconstructed profiles are compared to turbulence profiles estimated from meteorological data and a good agreement is found. We discuss CO-SLIDAR's contribution in the $C_n^2$ profilers' landscape and we propose some improvements of the instrument. ",
        "introduction": "The vertical distribution of atmospheric turbulence strength, known as the $C_n^2$ profile, is a key-point in the development of next-generation Adaptive Optics (AO) systems. Wide Field AO (WFAO) such as Laser Tomographic AO (LTAO), Multi-Conjugated AO (MCAO) and Multi-Object AO (MOAO) are based on a tomographic reconstruction of the atmospheric turbulence. They require a detailed knowledge of the $C_n^2$ profile for AO design and performance evaluation purposes~\\cite{lelouarn-a-00, fusco-a-01,costille-p-11} and for the forecast of turbulence conditions~\\cite{masciadri-a-13} to perform flexible scheduling of the observations with the instruments. $C_n^2$ profiles with a sub-kilometric vertical resolution in the $\\left[0 \\;; \\; 20\\right]$~km vertical range are necessary for these objectives~\\cite{costille-p-11}. \\newpage Many $C_n^2$ profilers have been developed through the last decades. SLODAR (SLOpe Detection And Ranging)~\\cite{wilson-a-02, butterley-a-06} uses correlations of slopes measured on a binary star with a Shack-Hartmann wavefront sensor (SH) to estimate the turbulent profile. Correlations of slopes allow a high sensitivity to ground and low altitude layers, but the performance is degraded at high altitudes. Other methods use correlations of scintillation. Among them SCIDAR (SCIntillation Detection And Ranging)~\\cite{rocca-a-74}, which works on a double star, and MASS (Multiple Aperture Scintillation Sensor)~\\cite{tokovinin-a-03b}, which uses a single star but only measures six turbulent layers. These techniques work better at high altitude and do not allow the estimation of near-ground turbulence, because they employ scintillation, whose variance is proportional to the altitude raised to the power $5/6^{\\mathrm{th}}$. The Generalized SCIDAR (G-SCIDAR)~\\cite{avila-a-97, fuchs-a-98} is an improvement of the SCIDAR that allows the method to be sensitive to the ground layer. The plane of the detector is the optical conjugate of a plane at a distance $h_{\\mathrm{gs}}$ which is below the telescope pupil, and this extra-propagation length enables G-SCIDAR to detect near-ground turbulence.  CO-SLIDAR~\\cite{vedrenne-a-07} is a recent $C_n^2$ profiler, combining sensitivity to both low and high altitude layers, which jointly uses correlations of slopes and of scintillation, both measured on a binary star with a SH. CO-SLIDAR's first validation in simulation have been presented in~\\cite{vedrenne-a-07}. The next step is a full experimental validation of the concept. To prepare the experiment, some experimental tests were performed in~\\cite{voyez-p-11} and an end-to-end simulation of the estimation of the $C_n^2$ profile in a practical case was detailed in~\\cite{voyez-p-12}.  This paper presents two improvements on the CO-SLIDAR reconstruction method and the first on-sky results of the CO-SLIDAR $C_n^2$ profiler. Detection noises are now taken into account and the resulting bias is subtracted in the reconstruction process. Error bars on the profile are now computed along with the estimated profile. The latest performance of CO-SLIDAR is first described in a simulated case. CO-SLIDAR is then tested on-sky, on the $1.5$~m MeO telescope. Images on a binary star are acquired to extract slope and scintillation data. Their correlations are computed so as to estimate $C_n^2$ profiles. The CO-SLIDAR profiles are estimated with their error bars. Results are compared to those obtained from correlations of slopes only or of scintillation only. Then, the CO-SLIDAR profiles are compared to profiles estimated from meteorological data, from NCEP (National Centers for Environmental Prediction)/NCAR (National Centers for Atmospheric Research) Reanalysis data base~\\cite{hach-a-12}. We finally discuss CO-SLIDAR's contribution as a new $C_n^2$ profiler and we suggest some improvements.  This paper is organized as follows. In Section~\\ref{sec:coslidar}, we recall CO-SLIDAR's principle and formalism and we describe the improvements on the reconstruction. In Section~\\ref{sec:simu} we show the latest performance of the CO-SLIDAR in a simulated case. Section~\\ref{sec:on_sky_results} is dedicated to the on-sky experiment and results. In Section~\\ref{sec:position} we discuss CO-SLIDAR's contribution in the $C_n^2$ profilers' landscape. Our conclusions and perspectives are given in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper, we have presented the first-on sky $C_n^2$ profiles estimated with the CO-SLIDAR profiler. The improved reconstruction method has been described and its latest performance has been illustrated in an end-to-end simulation, before testing it in a real astronomical observation on the $1.5$~m MeO telescope. $C_n^2$ profiles have been estimated with the CO-SLIDAR method and compared to those reconstructed from correlations of slopes only or of scintillation only, highlighting the fact that the two kinds of correlations are complementary and allow the reconstruction of low and high altitude layers. With this CO-SLIDAR instrument, we have reconstructed $30$ turbulent layers, with a resolution of $600$~m, from the ground up to $17$~km. We have compared our results with the $C_n^2$ profiles deduced from NCEP/NCAR Reanalysis, and we have found a good agreement between the estimations. In a nutshell, we have shown that CO-SLIDAR is a sub-kilometric resolution $C_n^2$ profiler working on meter class telescopes. The method now needs cross-calibrations with other optical profilers in multi-instrument campaigns. Considered future works include improvements to gain a better sky-coverage, performing real-time $C_n^2$ estimations, and extensions to wind and outer scale profiling."
    },
    "1402/1402.6545_arXiv.txt": {
        "abstract": "GRB 120326A is an unusual gamma-ray burst (GRB) which has a quite long plateau and a very late rebrightening both in X-ray and optical bands. The similar behavior of the optical and X-ray light curves suggests that they maybe have a common origin. The long plateau starts from several hundred seconds and ends at tens of thousands seconds. The peak time of the late rebrightening is about 30000 s. We analyze the energy injection model by means of numerical and analytical solutions, considering both the wind environment and ISM environment for GRB afterglows. We especially study the influence of the injection starting time, ending time, stellar wind density (or density of the circumburst environment), and injection luminosity on the shape of the afterglow light curves, respectively. We find that the light curve is largely affected by the parameters in the wind model. There is a ``bump'' at the late time only in the wind model too. In the wind case, it is interesting that the longer the energy injected, the more obvious the rebrightening will be. We also find the peak time of bump is determined by the stellar wind density. We use the late continuous injection model to interpret the unusual afterglow of GRB 120326A. The model can well fit the observational data, however, we find that the time scale of the injection must be larger than ten thousands seconds. This implies that the time scale of the central engine activity must be more than ten thousands seconds. This can give useful constraints on the central engine of GRBs. We consider a new born millisecond pulsar with strong magnetic field as the central engine. On the other hand, our results suggest that the circumburst environment of GRB 120326A is very likely a stellar wind. ",
        "introduction": "} Gamma-ray bursts (GRBs) are violent phenomena in the Universe, which radiate tremendous energy about $10^{51} - 10^{54}$ ergs from fractions of a second to tens of seconds. The widely accepted model is the fireball shock model (Goodman 1986; Rees \\& M\\'{e}sz\\'{a}ros 1992, 1994;  M\\'{e}sz\\'{a}ros \\& Rees 1992; Piran 1999). In this model, it is believed that the electrons are accelerated by internal shocks or external shocks. Internal shocks are likely developed near the site of optically thin fireball (Rees \\& M\\'{e}sz\\'{a}ros 1994), which give birth to the prompt emission due to collisions of shells with each other. Soon after the burst, the relativistic ejecta continue to spread out to form an external shock. The external shock sweeps up interstellar medium (Blandford \\& McKee 1976; Piran et al. 1993; Sari et al. 1998), which is believed to produce X-ray, optical/IR and radio emission, i.e. the afterglow.  The {\\it Swift} satellite (Gehrels et al. 2004), which was launched by NASA on November 20, 2004, opened a new era for GRBs researches. The X-ray Telescope (XRT) (Burrows et al. 2005b) onboard the {\\it Swift} has detected several hundreds of X-ray afterglows and the results are fruitful (M\\'{e}z\\'{a}ros 2006; Zhang 2007). In the summarization of the X-ray afterglow data, a canonical X-ray afterglow light curve includes five components (Zhang et al. 2006; Nousek et al. 2006). These five components include the steep decay phase (Zhang et al. 2007; Zhang et al. 2009), the shallow decay phase (Liang et al. 2007), the normal decay phase (Willingale et al. 2007), the post-jet break phase (Liang et al. 2008) and X-ray flares (Burrows et al. 2005a; Dai et al. 2006), respectively. For many bursts, we do not observe all components simply because of inadequate observations. Lots of samples only consist of single-power component (Liang et al. 2009). Recently, Zhang et al. (2013) studied the long-term central engine activities in the X-ray afterglow. Hou et al. (2014) studied a special sample, GRB 130925A, which appeared to a series of flares in the X-ray afterglow. Different samples show different temporal structures, increasing our perplexity. Meanwhile, Li et al. (2012) had systematically decomposed the optical afterglow light curves, indicating that the structures and composition of the optical afterglows were more complex than the X-ray afterglows.  For X-ray afterglows, the shallow decay phase is still a puzzle (Zhang 2007). The broadband afterglows, which usually decay as a power-law function of time with an index of $\\alpha \\sim 1.2$ (normal decay phase), are believed to be associated with the external shock. If the external shocks are refreshed by continuous energy injection into the blast wave, a shallow decay phase prior to the normal decay phase could be observed. The shallow decay may be due to the following mechanisms: (i) Energy injection invoking a long-term central engine (Dai \\& Lu 1998; Zhang \\& M\\'{e}sz\\'{a}ros 2001; Zhang et al. 2006; Nousek et al. 2006), (ii) Late internal shock model ( Zhang \\& M\\'{e}sz\\'{a}ros 2002; Zou et al. 2013), (iii) Off-axis jet model (Eichler \\& Granot 2006; Toma et al. 2006), (iv) Central engine model (Kumar et al. 2008; Geng et al. 2013 ). To determine which one of these models is correct, we need more observational data.  A quite long plateau followed by a very late rebrightening is observed in the afterglow of GRB 120326A in X-ray and optical bands, which is observed for the first time. The feature of rebrightening behavior around 30000 s is not easy to understand. The phenomenon is so special that a lot of telescopes observe this burst later. We collect a lot of observational data, including early and late afterglow in optical band. It can potentially helps us to understand the underlying mechanism of this afterglow. Unlike the common X-ray flares or the prompt pulses, which are usually characterized by a rapid rise and an exponent decay (Kocevski et al. 2007; Norris et al. 2005; Chincarini et al.2007, 2010; Margutti et al. 2011; Li et al. 2012), GRB 120326A shows a slight bump overlapping the afterglow light curve.  In this work, we mainly use the energy injection model in the wind case to explain the particular phenomenon of GRB 120326A. The plan of the paper is the following: in Section 2, we briefly describe the observation data. We review the energy injection model and fit the X-ray data in Section 3. In Section 4, we present our conclusions and discussion. The cosmological parameters $H_{0}$ = 71 km $\\rm s^{-1}$ $\\rm Mpc^{-1}$, $\\Omega_{M}$ =0.3, and $\\Omega_{\\Lambda}$ = 0.7 have been adopted throughout our study. ",
        "conclusions": "GRB 120326A is an unusual GRB, which shows a quite long plateau and a very late rebrightening in X-ray and optical bands at the same time. The similar shape in X-ray and optical bands of GRB 120326A suggests the same origin for them. The long plateau starts at several hundred seconds and ends at tens of thousands seconds. The peak time of the late rebrightening is about 30000 s. Because of such a curious phenomenon, a lot of telescopes did the follow-up observations of this GRB. We collect a lot of observational data, having early and late afterglow data to deepen our understanding of the energy injection mechanism.  Plateaus and rebrightenings/bumps in GRB afterglow light curves usually hint late reactivity of the GRB central engine (Dai \\& Lu 1998; Zhang \\& M\\'{e}sz\\'{a}ros 2001; Geng et al. 2013). However, it is rare that two phenomena take place in one GRB. In this study, we consider a new born millisecond pulsar with strong magnetic field as the central engine (Dai \\& Liu 2012). Zhang and Yan (2011) also suggested that the central engine could produce a Poynting-flux dominated outflow. Due to the strong magnetic field and rapid rotation, the new born millisecond pulsar can release the energy comparable to that of the impulsive energy of the initial fireball. In this case, the very early afterglow may be directly powered by ultra-relativistic mildly magnetized outflows (Fan et al. 2004). It is also very possible that the continuous energy release by the pulsar will be injected into the external shock driven by the initial jet, which is currently the standard scenario for afterglow plateaus.  We analyze the energy injection model by means of numerical and analytical solutions, considering both the wind scenario and ISM scenario. The influence of the injection starting time, ending time, stellar wind density (or density of the circumburst environment), and injection luminosity on the shape of the afterglow light curves are explored. We find that the light curve of the afterglow is largely affected in the wind model by the parameters. We also find that although the light curve is slightly affected in the ISM case by these parameters, there is a significant ``bump'' at the late time only in the wind model.  For different GRBs, the circumburst environment (wind parameter $A_{*}$), the luminosity ($L_0$) and duration ($T_{\\rm start}$, $T_{\\rm end}$) of the energy injection may vary markedly. This may be one possible explanation to the observed diverse afterglow light curves. From Figure 3a and 4a, we show that the plateau and the late bump is an indicator of a wind type circumburst environment. GRB 120326A with such temporal behavior is an ideal template. Our results suggest that the circumburst environment of GRB 120326A is very likely a stellar wind. With reasonable parameter values, we give a good fitting to the X-ray and optical afterglow light curves of GRB 120326A (see Fig. 2). The model parameters from our fits are $A_{*}=0.45$, $T_{\\rm start}=600$ s, $T_{\\rm end}=3\\times10^4$ s, $L_0=1.70\\times10^{49}$ erg $\\rm s^{-1}$."
    },
    "1402/1402.1223_arXiv.txt": {
        "abstract": "We report on balanced SIS receivers covering the astronomical important 180$-$720 GHz submillimeter atmospheric window. To facilitate remote observations and automated spectral line surveys, fully synthesized local oscillators are employed. High-current-density Nb-AlN-Nb superconducting-insulating-superconducting (SIS) tunnel junctions are used as the mixing element. The measured double-sideband (DSB) 230 GHz receiver noise temperature, uncorrected for optics loss, ranges from 50~K at 185~GHz, 33~K at 246~GHz, to 51~K at 280~GHz.  In this frequency range the mixer has a DSB conversion gain of 0~$\\pm$~1.5~dB. The measured 460~GHz double-sideband receiver noise temperature, uncorrected for optics loss, is 32~K at 400~GHz, 34~K at 460~GHz, and 61~K at 520~GHz. Similar to the 230 GHz balanced mixer, the DSB mixer conversion gain is 1~$\\pm$~1~dB.  To help optimize performance, the mixer IF circuits and bias injection are entirely planar by design. Dual-frequency observation, by means of separating the incoming circular polarized electric field into two orthogonal components, is another important mode of operation offered by the new facility instrumentation. Instrumental stability is excellent supporting the LO noise cancellation properties of the balanced mixer configuration. In the spring of 2012 the dual-frequency 230/460 SIS receiver was successfully installed at Caltech Submillimeter Observatory (CSO), Mauna Kea, HI. ",
        "introduction": "\\label{sect:intro} \\IEEEPARstart{T}he Caltech Submillimeter Observatory (CSO) is located on top of Mauna Kea, Hawaii, at an altitude of 4.2~km. To facilitate deep integrations, stable baselines, and automated high resolution spectral line surveys significant efforts \\cite{kooi_2012, kooi_2007} have been expanded to develop a family of synthesized, remote programmable dual-color balanced SIS receivers. To this effect four tunerless balanced-input waveguide receivers have been developed to cover the important 180$-$720~GHz frequency range \\cite{spie2002, spie2004}. The new heterodyne facility instrumentation allows dual-frequency (two-color) observations in the 230/460~GHz and 345/660~GHz atmospheric windows. The 230/460 balanced receiver system has been installed and is operational at the observatory since May 2012. Unfortunately, deployment of the 345/650 balanced receiver(s) is presently on hold due to funding difficulties.  Dual-frequency observation is an important mode of operation offered by the new facility instrumentation. Dual band observations are accomplished by separating the horizontal (H) and vertical (V) polarizations of the incoming signal and routing them via folded optics to the appropriate polarization sensitive balanced mixer. Scientifically this observation mode facilitates pointing for the higher receiver band under mediocre weather conditions and a doubling of scientific throughput under good weather conditions.  Balanced configurations were chosen for their inherent local oscillator (LO) spurious tone and amplitude (AM) noise cancellation properties. (It was also judged to be an optimal compromise between scientific merit and finite funding). Unique to the CSO, wide RF bandwidth is favored \\cite{kooi_2007}, allowing the same science to be done with fewer instruments. In all the upgrade covers ALMA bands~5$-$9.  In principle, the balanced receiver configuration has the advantage that common mode amplitude noise in the local oscillator system is canceled, while at the same time utilizing all available LO power.  Both of these features afford automation over unprecedented wide RF bandwidth, covering the entire 180$-$720~GHz submillimeter atmospheric windows with just four LOs.  Receiver noise temperatures and in situ measured instrumental Allan Variance stability times are excellent and are consistent with the use of balanced receiver technology.  The SIS junctions are capable of a 13~GHz bandwidth, though due to band limiting isolators and low noise amplifiers the operational IF bandwidth of the CSO receivers is presently 4$-$8~GHz.  To maximize the RF bandwidth, we explore the use of high-current-density AlN-barrier SIS technology combined with a broad bandwidth full-height waveguide to thin-film microstrip transition \\cite{kooi_probe}. Compared to AlO$_x$-barriers, advantages of AlN tunnel barriers are a low $\\omega$RC product (increased RF bandwidth) and enhanced chemical robustness. Even if optimal RF bandwidth is not a requirement, a low $\\omega$RC product provides a more homogeneous frequency response and increased tolerance to errors in device fabrication.  To process the required IF bandwidth, the CSO has acquired a Fast Fourier Transformer Spectrometer (FFTS) from Omnisys Instruments, Sweden. \\footnote{ Omnisys Instruments AB., August Barks gata 6B, SE-421 32 V\\\"{a}stra Fr\\\"{o}lunda SWEDEN. [Online] Available: http://www.omnisys.se} This spectrometer facilitates 8 GHz of processing bandwidth with a resolution of 268 KHz/channel, or 3724 channels/GHz. The 8 GHz Omnisys FFTS comes in a 19\" rack and has two built-in IF processor modules (4$-$8 GHz each), an embedded controller module, a synchronization module, and power supply. ",
        "conclusions": "To facilitate deep integrations, automated line surveys, and remote observations a suite of fully synthesized dual-color balanced receivers covering the 180$-$720 GHz submillimeter atmospheric frequency range (ALMA B5$-$B9) have been developed for the Caltech submillimeter Observatory. It was judged an optimal compromise between scientific merit and finite funding. In the development wide RF bandwidth was favored \\cite{kooi_2007, kooi_2012}, allowing the same science to be done with fewer instruments.   High-current-density (25~kA/cm$^2$) AlN-barrier SIS technology has been used, facilitating the very wide instantaneous RF bandwidth presented. For the 230~GHz balanced receiver we obtain in the 180$-$280~GHz frequency range a receiver sensitivity of 33$-$50~K DSB. The mixer conversion gain is relatively constant at 0~$\\pm$~1 dB. The 460~GHz balanced receiver also has a noise temperature in the range 40$-$50~K DSB, with a mixer conversion gain of 1~$\\pm$~1 dB. In the case of the 460~GHz receiver the very low receiver noise temperature is understood to be the result if four factors: A cooled last stage multiplier, $\\sim$12~dB of AM noise rejection, moderate mixer conversion gain, and high optical throughput. Conversely, the 230 GHz receiver noise temperature suffers from the fast optics required to achieve proper telescope illumination and less of a reduction in LO noise as is the case with the 460~GHz instrument.  The amplitude noise rejection is 12~dB $\\pm$~3~dB for both receivers. This result slightly exceeds the theoretical estimate of \\cite{kooi_2012}. Instrument stability, as measured in situ at the telescope, is excellent and argues for the development of submillimeter and terahertz balanced receiver configurations. The presented results are confirmed by actual observations and the obtained high quality spectral baselines.  Unfortunately, deployment of the 345/650 balanced receiver(s) is presently on hold due to funding difficulties."
    },
    "1402/1402.5145.txt": {
        "abstract": "{ Galaxy surveys targeting emission lines are characterising the evolution of star-forming galaxies, but there is still little theoretical progress in modelling their physical properties. We predict nebular emission from star-forming galaxies within a cosmological galaxy formation model. Emission lines are computed by combining the semi-analytical model \\sag\\ with the photoionisation code \\mapp. We characterise the interstellar medium (ISM) of galaxies by relating the ionisation parameter of gas in galaxies to their cold gas metallicity, obtaining a reasonable agreement with the observed \\ha, \\oii, \\oiii\\ luminosity functions, and the the BPT diagram for local star-forming galaxies. The average ionisation parameter is found to increase towards low star-formation rates and high redshifts, consistent with recent observational results. The predicted link between different emission lines and their associated star-formation rates is studied by presenting scaling relations to relate them. Our model predicts that emission line galaxies have modest clustering bias, and thus reside in dark matter haloes of masses below $M_{\\rm halo} \\lesssim 10^{12} {[\\rm h^{-1} M_{\\odot}]}$. Finally, we exploit our modelling technique to predict galaxy number counts up to $z\\sim 10$ by targeting far-infrared (FIR) emission lines detectable with submillimetre facilities.} ",
        "introduction": "Emission lines are a common feature in the spectra of galaxies. A number of physical processes occurring in the interstellar medium (ISM) of galaxies can be responsible for their production, such as the recombination of ionised gas, collisional excitation or fluorescent excitation \\citep{osterbrock89,stasinska07}. Their detection allows the exploration of the physical properties of the medium from which they are produced. Active galactic nuclei (AGN), for instance, can photoionise and shock-heat the gas in the narrow-line region surrounding the accretion disks of supermassive black holes \\citep{groves04a,groves04b}. Also, cold gas falling into the galaxy's potential well might radiate its energy away via collisional excitation \\citep{dijkstra09}.  Likewise, star-forming activity in galaxies can be traced by the detection of emission lines \\citep{kennicutt83,calzetti12a}. Young, massive stars produce photoionising radiation which is absorbed by the neutral gas of the ISM leading to the production of emission lines. Given the short life-times of massive stars (of the order of a few Myr), emission lines indicate the occurrence of recent episodes of star formation. On the contrary, continuum-based tracers of star forming (SF) activity, such as the UV continuum, are tracing the star formation rate (SFR) over a much larger time-scale of the order of $\\sim$100 Myr. Regardless of the technique used, mapping the cosmic evolution of star formation to understand the formation and evolution of galaxies has become an important challenge in extragalactic astrophysics { \\citep{madau98,hopkins04,madau14}}.  The link between emission line fluxes and the properties of the gas in the ISM have naturally made their detection a common tool in extragalactic astrophysics. Theoretical modelling of emission lines has focused mostly in providing a tool to learn about the dynamics and structure of the gas in galaxies. Typically, the interpretation of line fluxes and line ratios is done by making use of a photoionisation radiative transfer code \\citep{ferland98,dopita00,kewley01}. Line ratios and fluxes can then be translated into gas metallicities, temperatures, densities and the ionising structure of the gas \\citep[e.g.][]{kewley08,levesque10}. By using a number of ad-hoc assumptions, it is also possible to include nebular emission into a stellar population synthesis code \\citep{charlot01,panuzzo03}. Individual objects can be studied in great detail when both optical and infrared emission lines are available, constraining the properties of the multi-phase ISM and the contribution of each ISM component to the total emission line budget \\citep[e.g.][]{abel05,cormier12}.  Analysis of the emission line budget has also been done in large galaxy datasets. In the local spectroscopic sample of galaxies from the Sloan Digital Sky Survey (SDSS), \\citet{tremonti04} made use of several line ratios finding a global relation between the gas metallicity and the stellar mass of galaxies. The exact shape and normalisation of this relation is somewhat controversial, since metallicity diagnostics using different line ratios can lead to different values of metallicity \\citep{kewley08,cullen13}. Despite the uncertainty in the exact value of the gas metallicity inferred by measuring line ratios, it is clear that the mass-metallicity relation encodes information about the evolution of the chemical enrichment of galaxies, which in turn depends on the formation history of galaxies, their stellar mass growth, the feedback mechanisms regulating the star-formation processes, and the merging histories of galaxies.  The narrow-band technique to find emission lines in high redshift galaxies has been extensively applied to large samples at different redshifts. The HiZELS survey \\citep{geach08} has produced a sample of high redshift star-forming galaxies large enough to measure the luminosity function and clustering of \\ha\\ and \\oii\\ emitters over the redshift range $0.4<z<2.2$ \\citep{sobral10,geach12,hayashi13,sobral13}. More recently, using slitless grism spectroscopy from the Hubble Space Telescope, the PEARS survey sampled a large number of  \\ha, \\oii\\ and \\oiii\\ emission line galaxies (ELGs) over the redshift range $0<z<1.5$ \\citep{xia11,pirzkal13}.  At high redshifts, hydrogen recombination and forbidden lines are also targeted to derive galaxy redshifts and also physical properties, such as the SFR and the gas metallicitiy { \\citep[e.g.][]{maiolino08,maier06}}. Galaxies at high redshifts have been found to have lower gas metallicities for a given stellar mass compared to their $z=0$ counterparts, thus suggesting an evolution of the mass-metallicity relation \\citep{erb06,maiolino08,yuan13}. More recently, a fundamental relation between gas metallicity, stellar mass and star-formation rate has been suggested \\citep{mannucci10,lara-lopez10} which seems to hold independent of redshift \\citep[see, however,][]{troncoso14}.  From a cosmological perspective, the search for the most distant galaxies is key to understand the sources that re-ionised the Universe \\citep{meiksin09,bunker10,raicevic11}. Searching for the Lyman break {\\citep{steidel96}} is a common technique to select galaxy candidates for these extreme high redshift galaxies. Normally, spectroscopic confirmation of their redshift is needed, which consists in detecting a strong emission line, such as \\lya\\ \\citep[e.g.][]{pentericci07,stark10,schenker12}. Emission lines are also targeted in large area surveys to measure the large scale structure of the Universe at high redshifts. Cosmological probes such as the detection of baryonic acoustic oscillations (BAOs) and the growth rate of cosmic structures at different redshifts are powerful tests that could revolutionise our cosmological paradigm. Ambitious programs have been proposed, including space missions such as EUCLID, whose goal is to detect $\\sim 10^8$ \\ha\\ emitters in the redshift range $0.5<z<2.2$, \\citep{laurejis10}, ground-based programs such as HETDEX \\citep{hill08,blanc11}, which expects to detect $\\sim 10^6$ \\lya\\ emitters over the redshift range $2.5<z<3.5$, and multi narrow-band wide field surveys such as PAU and J-PAS which have also the potential to study the large scale structure of the Universe traced by emission line galaxies { \\citep{abramo12,benitez14}}.  However, despite the significant observational progress which has increased the depth and area in surveys of ELGs, little theoretical work has been made aiming at understanding the emission line galaxy population as a whole. Emission from the $^{12}$CO molecule has been studied by \\citet{lagos12} in a  galaxy formation model, where the molecular gas content is key to predict the luminosity from different rotational transitions of $^{12}$CO. Also, a number of studies have focused in the \\lya\\ line, which, unlike the lines generated by the $^{12}$CO molecule, is originated in HII regions \\citep[e.g.][]{ledelliou06,dayal10a,kobayashi10,orsi12}. The intrinsic production of \\lya\\ photons from recombination of photoionised hydrogen is not a challenge to the models, but the high \\lya\\ cross section of scattering with hydrogen atoms results in the path lengths of photons being large enough to make them easily absorbed by even small amounts of dust. Predictions for \\lya\\ emitters were presented in detail in \\citet{orsi12} including the escape of \\lya\\ photons through galactic outflows computed with a Monte Carlo radiative transfer model.  In this work we extend the modelling of ELGs by focusing in other nebular emission lines in the optical and far-infrared (FIR) range. The ELG population is studied by making use of a fully fledged semi-analytical model of galaxy formation. The \\sag\\ model \\citep[acronym for Semi-Analytic galaxies;][]{cora06,lagos08} makes use of N-body dark matter simulations to follow the abundances and merging history of dark matter haloes where astrophysical processes shape the formation and evolution of galaxies. Predictions of \\ha\\ emission in star-forming galaxies were presented in \\citet{orsi10} making use of the semi-analytical model \\galform. Here we extend the range of predictions by combining the \\sag\\ model with the photoionising code \\mapp\\ \\citep{dopita00,levesque10} to predict the luminosities of forbidden lines in the optical such as \\oii\\ and \\oiii\\ and far-infrared (FIR) lines such as \\nii.  The paper is organised as follows: Section \\eqref{sec.models} briefly describes the semi-analytical model \\sag\\ and the photoionisation code \\mapp\\ that we use in this work. Section \\eqref{sec.elgmodel} presents our strategy to compute the emission lines of galaxies. Section \\eqref{sec.results} shows how our predictions compare to observational measurements and describes predictions of the properties of ELGs. In Section \\eqref{FIR} we present predictions for very high redshift galaxies that could be observed with submillimetre facilities. Then, in Section \\eqref{sec.discussion} we discuss the validity of our modelling strategy. Finally, in Section \\eqref{sec.conclusions} we summarise and present the main conclusions of this work. ",
        "conclusions": "\\label{sec.conclusions}  Throughout this paper we explore the properties and evolution of emission line galaxies within a hierarchical cosmological model. This is the first attempt to model in detail the strong emission lines of galaxies in the optical and FIR range and to study their statistical properties as a galaxy population.  Our method consists in combining the semi-analytical model of galaxy formation \\sag\\ with the photoionisation and shock code \\mapp\\ to predict the emission line luminosities of galaxies. We make the critical assumption of relating the ionisation parameter of galaxies with the cold gas metallicity through Eq. \\eqref{eq.qZ}. We choose the constant values in this relation to achieve a reasonable agreement of the BPT diagnostic and the \\oii\\ and \\oiii\\ LFs with the observational data available. We take advantage of the power of semi-analytical models to explore the properties of the ELG population up to redshifts $z=10$, and present predictions of the detectability of FIR lines that can be targeted with millimetre and submillimetre facilities to explore the high redshift Universe.  We summarise our main findings as follows:  \\begin{itemize}  \\item Overall, our model predicts ELG luminosity functions in close agreement with current observational data. The model predicts an overabundance of \\ha\\ emitters at $z \\lesssim 0.4$, but is reasonably consistent with the observed LFs in the range $0.9<z<2.2$. The luminosity functions of \\oiii\\ are find to be in partial agreement with observations. The abundance of \\oii\\ emitters at high redshifts is remarkably close to observed values, { but it is inconsistent with the observed abundance at $z=0$}. There are many reasons that can explain the disagreements between model predictions and observational measurements. First, our model relating the ionisation parameter $q$ with the cold gas metallicity $Z$ could be an oversimplification of the complex internal gas dynamics in HII regions of star-forming galaxies. Secondly, the ionisation parameter our model computes depends solely on the cold gas metallicity, which implies that if the chemical enrichment histories of galaxies are not realistic in our model then the resulting ionisation parameter values we obtain for high redshift galaxies would be incorrect. This would be reflected in the luminosity function of \\oii\\ and \\oiii. Finally, contamination from AGNs, shock excitation or LINERs might affect the observed luminosity functions in a rather complicated way. \\citet{kewley13a} shows that the strong line ratios of star-forming galaxies and also AGNs are expected to evolve with redshift, and that diagnostics such as the traditional BPT diagram cannot be applied to discriminate the source of emission lines in the same way that this is done to local galaxies.  \\item Our model predicts that star-forming galaxies at high redshifts have higher ionisation parameters than their counterparts at $z=0$. A similar conclusion has been suggested in a number of observations \\citep[e.g.][]{brinchmann08,hainline09,nakajima13,richardson13}. This is an expected result, since galaxies accumulate metals throughout their star-forming history, and so the bulk of the galaxy population at high redshifts is predicted to have less metal abundances than local galaxies. Since the ionisation parameter in our model scales inversely with the metallicity, high redshift galaxies have therefore higher values of the ionisation parameter.  \\item We find that line luminosities can be used to infer the star-formation rate of galaxies, in spite of the scatter of the relation. We fit a polynomial functional form which depends on the emission line luminosity and the redshift of the galaxies. The scaling relations we present are in general accurate to within an order of magnitude for low luminosities and their accuracy improves towards brighter line luminosities. The less favoured emission line to infer the SFR from is found to be \\oiii.  \\item We find that our model predicts modest values for the clustering of emission line galaxies, roughly consistent with observational estimates of \\ha\\ clustering. The bias of emission line galaxies is found to increase towards brighter galaxy samples, especially in high redshift emission line galaxies.\\oiii\\ emitters present a strong connection between the bias and the line luminosity, and this correlation is found to be weaker in \\oii\\ and \\ha\\ emitters. Overall, ELGs are found to be hosted by dark matter haloes of mass $M_{\\rm halo} \\lesssim 10^{12} {\\rm [h^{-1}M_{\\odot}]}$.  \\item We explore the detectability of FIR emission lines with submillimetre observations. \\oiiiee\\ is found to be the strongest line, although we do not explore the more common target \\cii. A dedicated survey with ALMA covering several square arc minutes is predicted to find a handful of \\oiiiee\\ emitters at $z=10$ and \\nii\\ emitters at $z=7$ if the flux density probes down to a depth of $0.01 {\\rm [mJy]}$.  \\end{itemize}  This is our first paper of a series exploring the properties of strong optical and FIR nebular emission in galaxies. Several improvements can be made to the model to account for more physical processes. For instance, the mid-plane pressure of the ISM could be used to infer an effective ionisation parameter. A detailed study of the chemical enrichment of galaxies throughout cosmic time could also prove to be key in validating or ruling out our simple modelling of the ionisation parameter.  Regardless of the limitations, this work constitutes the first attempt to explore the properties of emission line galaxies in a cosmological setting. With the advent of ongoing and future large area surveys of emission line galaxies, a detailed understanding of this galaxy population will be regarded as crucial to extract information about the large scale structure of the Universe, and the process of galaxy evolution."
    },
    "1402/1402.1824_arXiv.txt": {
        "abstract": "We present new Very Large Array 6$\\,$cm H$_2$CO observations toward four extragalactic radio continuum sources (B0212+735, 3C$\\,$111, NRAO$\\,$150, BL$\\,$Lac) to explore the structure of foreground Galactic clouds as revealed by absorption variability. This project adds a new epoch in the monitoring observations of the sources reported by Marscher and collaborators in the mid 1990's. Our new observations confirm the monotonic increase in H$_2$CO absorption strength toward NRAO$\\,$150. We do not detect significant variability of our 2009 spectra with respect to the 1994 spectra of 3C111, B0212+735 and BL$\\,$Lac; however we find significant variability of the 3C111 2009 spectrum with respect to archive observations conducted in 1991 and 1992. Our analysis supports that changes in absorption lines could be caused by chemical and/or geometrical gradients in the foreground clouds, and not necessarily by small scale ($\\sim 10\\,$AU) high density molecular clumps within the clouds. ",
        "introduction": "The 6$\\,$cm line of formaldehyde (H$_2$CO) is a powerful probe to study molecular gas in the Milky Way and in other galaxies (e.g., Mangum et al. 2008; Dickel et al. 2001). Soon after its discovery (Snyder et al. 1969), the H$_2$CO 6$\\,$cm transition was detected in absorption toward radio continuum sources as well as against the Cosmic Microwave Background (e.g., Palmer et al. 1969; Dieter 1973). Now, more than 40 years of observations have shown that the 6$\\,$cm H$_2$CO line can trace large ($>5$pc) molecular clouds (e.g., Araya et al. 2007). Such large structures are not homogeneous, for example, Dieter-Conklin (2010) concluded that molecular clouds traced by H$_2$CO absorption show different velocity dispersion when observed at effective resolutions greater or smaller than $\\sim 0.5\\,$pc, which gives direct evidence for sub-parsec structure traced by H$_2$CO.  Even smaller size-scales can be probed by H$_2$CO observations of absorption against compact extragalactic radio sources. In particular, Marscher et al. (1993) and Moore \\& Marscher (1995) used the VLA to monitor 6$\\,$cm H$_2$CO absorption toward four extragalactic radio sources during a period of $3.4\\,$years, and found evidence for H$_2$CO line variability toward three continuum sources (3C$\\,$111, NRAO$\\,$150, BL Lac; they found no significant change in the B0212+735 line profile based on observations conducted less than 2 years apart). As recognized by Marscher et al. (1993) and Moore \\& Marscher (1995), the variability most likely traces changes in our line of sight that occurs because we move with respect to the distant source and intervening cloud, i.e., the combined effect of Earth's revolution around the Sun and the cloud's internal kinematics and proper motion (see Dieter-Conklin 2009, for a discussion of this ``searchlight effect''). Marscher et al. (1993) estimated that the change in line of sight through the clouds enables probing physical scales of the order of $4\\,$AU per year, i.e., $\\sim 14\\,$AU for sources with 3.4$\\,$years monitoring data.  The use of 6$\\,$cm H$_2$CO absorption as probe for substructure (as revealed by spectral line variability) has clear advantages, in particular, the supercooling effect of the lowest K-doublet transition due to H$_2$ - H$_2$CO collisions  (Townes \\& Cheung 1969; Thaddeus 1972; Garrison et al. 1975; Troscompt et al. 2009a) enhances the optical depth against continuum sources. Also, the separation between the main group of five hyperfine lines (centered near the $F = 2-2$ component) and the $F = 1-0$ component is approximately 18.2$\\,$kHz (equivalent to 1.13\\kms; see Tucker et al. 1970) and therefore the $F = 1-0$ line does not significantly overlaps with the main hyperfine components in diffuse/translucent regions and dark clouds (which have linewidths of $\\lesssim 1$\\kms; see for example Araya et al. 2006). Thus, if variability is detected in the main absorption feature, variability of the hyperfine component is also expected, which is a useful test to assess whether artifacts are present in the data.  The monitoring of the H$_2$CO line profiles started by Marscher and collaborators did not continue. In this paper we report new observations of the sources observed by Moore \\& Marscher (1995) to investigate longer term variability in the absorption profiles and the implication of the variability on the structure of the clouds. ",
        "conclusions": "\\subsection{Reproducibility of Spectral Line Results}  To check the effect of different calibration procedures, we reduced all archive observations of the sources which have been already published by Marscher et al. (1993) and Moore \\& Marscher (1995). We applied completely independent flagging, calibration and imaging procedures. In our analysis of the physical parameters of the clouds, the peak optical depths and linewidths are input parameters (see section 4.2), while Moore \\& Marscher (1995) used equivalent widths defined as $\\int \\tau dv$. To compare our results to those of Moore \\& Marscher (1995), we measured the equivalent widths of our reduction of the M94 data. Our equivalent widths agree with Moore \\& Marscher (1995) values within statistical uncertainties from the fit in all sources, which shows that differences in calibration procedures are not a significant source of error in determining variability. For example, in the case of the M94 observations of B0212+735, Moore \\& Marscher (1995) reported $\\int \\tau dv$ = 0.0893, $V_{LSR, F=2-2} = 3.56$\\kms, and $FWHM_{F=2-2} = 0.84$\\kms, while we measured $\\int \\tau dv = 0.091 \\pm 0.012$, $V_{LSR, F=2-2} = 3.57 \\pm 0.01$\\kms, and $FWHM_{F=2-2} = 0.81 \\pm 0.04$\\kms.  We detect a strong absorption line at $-17.2$\\kms~toward NRAO$\\,$150 (see Table~3, Figure~3). This line was included in the bandpass (different IFs) of the J91 and D90 epochs, however, the line was not included in the bandpass of the M94 and D92 epochs. As clearly seen in Figure~3, the line profile of the $-17.2$\\kms~component is perfectly well reproduced by our 2009 observations with respect to the previous epochs, showing that the variability of the $-10.5$\\kms~component in NRAO$\\,$150 is not an artifact. Likewise, the $-2.3$\\kms~velocity component in 3C111 shows no variability (compare red and black spectra in Figure~5) while the main absorption feature at $-1.0$\\kms~does.  Careful inspection of the B0212+735 and NRAO$\\,$150 spectra (Figures~1 and 3) shows that the bandpasses are not perfectly flat; specifically, the B0212+735 spectrum has a weak absorption feature at $\\sim$0\\kms~and NRAO$\\,$150 presents weak absorption features at $-$14 and $-$4.5\\kms~(highlighted with arrows). Absorption features at these velocities have been seen before in other molecular transitions. The CO, OH, and HCO$^+$ spectra of B0212+735 reported by Liszt \\& Pety (2012) show a velocity component at $\\sim 0$\\kms~in addition to the main component at $\\sim 3.5$\\kms. In the case of NRAO$\\,$150, molecular absorption features at $-$14 and $-$4.5\\kms~(in addition to the components at $-17.3$ and $-10.5$ and $-8.5$\\kms) have been reported by Pety et al. (2008) in CO, HCO$^+$, C$_2$H, and HNC; of all, the HNC spectrum is the most similar to the H$_2$CO spectrum. The same weak absorption features are seen at the other epochs (see Figure~3, left spectra). The remarkable reproducibility of these weak features demonstrates the superb stability of the VLA for long-term variability studies of spectral lines.   \\subsection{Physical Parameters}  We used the {\\tt Molpop-CEP} code (Elitzur \\& Asensio Ramos 2006) to explore the physical conditions of the foreground clouds based on the observed H$_2$CO absorption lines. We used the H$_2$CO -- para-H$_2$ collision rates of Troscompt et al. (2009b), and included He collisions using the rates by Green (1991)\\footnote{Electron collisions could be an important factor in the H$_2$CO excitation in translucent/diffuse clouds (e.g., Turner 1993), however development of a model including electron collisions is beyond the scope of this work.}. Diffuse/translucent clouds have densities between $\\sim 10^2$ and $\\sim 10^4\\,$cm$^{-3}$ (e.g., van Dishoeck \\& Black 1988), and according to Moore \\& Marscher (1995) may contain even higher density clumps [$\\sim 10^6\\,$cm$^{-3}$; however see Liszt \\& Lucas (2000) and our discussion below]. The H$_2$CO abundance relative to hydrogen may be somewhere between $10^{-9}$ and $\\sim 5 \\times 10^{-8}$; for example, Zhou et al. (1990) estimated an abundance of $2 \\times 10^{-9}$ toward B335; Turner (1993) report an abundance of $6 \\times 10^{-9}$; Troscompt et al. (2009a) obtained an abundance of $6 \\times 10^{-8}$. We ran grids of models at two kinetic temperatures (10 and 30$\\,$K), in a range of densities between 10$^2$ to 10$^6\\,$cm$^{-3}$ and H$_2$CO abundance ratios between $10^{-9}$ and $5 \\times 10^{-8}$, however, for the analysis, we further constrain the parameters as follows:  {\\it Temperature:} The extinction toward the extragalactic sources discussed here is $A_V \\sim 1\\,$mag or greater (e.g., Marscher et al. 1993; Schlegel et al. 1998; Pety et al. 2008; Schlafly \\& Finkbeiner 2011; Liszt \\& Pety 2012). Hence, the foreground clouds are in the diffuse/translucent boundary; which are expected to be hotter than $\\sim$10$\\,$K at their centers and as hot as $\\sim 50\\,$K at their boundaries (e.g., van Dishoeck \\& Black 1988). Given that the H$_2$CO molecules are not expected at the boundaries of the clouds but at higher depths where they may be shielded from the interstellar radiation field, a temperature of $\\sim 30\\,$K is assumed.  {\\it Density:} observations of 6$\\,$cm H$_2$CO absorption profiles in a number of environments including dark nebulae, pre-protostellar cores and low mass star forming regions with absorption against the CMB (e.g., Palmer et al. 1969; Heiles 1973; Dieter 1973; Young et al. 2004; Araya et al. 2006) have resulted in excitation temperature measurements of the 6$\\,$cm K-doublet of $\\lesssim 2\\,$K, which according to our model imply densities below $\\sim 10^5\\,$cm$^{-3}$. This is the well known result that the 6$\\,$cm H$_2$CO K-doublet thermalizes at densities of $\\sim 10^6\\,$cm$^{-3}$ (e.g., Troscompt et al. 2009a; Mundy et al. 1987). Thus, most of the H$_2$CO gas likely traces regions of densities below $\\sim 10^5\\,$cm$^{-3}$. A lower limit of the density can be inferred from our data as well: as highlighted in Figures~1 and 3 upper-right panels and discussed above, we detected some weak H$_2$CO absorption features. These velocity components are quite prominent in CO and HCO$^+$ (e.g., Pety et al. 2008), and may trace molecular gas at densities of $\\sim 10^2\\,$cm$^{-3}$ (e.g., Liszt \\& Lucas 2000; Pety et al. 2008). H$_2$CO is likely weaker because higher density conditions would be needed for stronger H$_2$CO absorption features (Pety et al. 2008). Thus, the H$_2$CO gas likely traces material at densities between $\\sim 10^3$ to $10^5\\,$cm$^{-3}$, i.e., higher than those traced by HCO$^+$ and lower than termalization densities; we therefore assume a density of $10^4\\,$cm$^{-3}$ for the analysis.  {\\it Abundance:} we simply assume an abundance of $5\\times10^{-9}$, i.e., similar to the abundance reported by Turner (1993) for high-latitude cirrus clouds.   Based on these assumptions, we can calculate the H$_2$CO and H$_2$ column densities, and the thickness ($\\Delta L$) of the absorbing region (assuming a slab geometry), i.e., the depth of the cloud that would result in the observed absorption spectrum for a given density, temperature and abundance. We note that the slab thickness should be considered an effective thickness, more precisely, a lower limit of the physical path which would differ from the true thickness of the cloud depending on the number of molecular clumps and their separation along the line of sight (see Marscher et al. 1993).  In the case of the H$_2$CO absorption region toward the core of 3C111 (see Figure~4), we obtained the following parameters and uncertainty ranges: $\\Delta L = 1.0 \\times 10^4\\,$AU ([$7 \\times 10^2$, $5 \\times 10^5$] AU), $N(H_2) = 1.5 \\times 10^{21}\\,$cm$^{-2}$ ([$1 \\times 10^{20}$, $5 \\times 10^{22}$] cm$^{-2}$), and $N(H_2CO) = 7.6 \\times 10^{12}\\,$cm$^{-2}$ ([$7 \\times 10^{12}$, $5 \\times 10^{13}$] cm$^{-2}$), where the ranges shown in square brackets were obtained by assuming densities between $10^3$ and $10^5\\,$cm$^{-3}$ and formaldehyde abundances with respect to hydrogen between $10^{-9}$ and $5 \\times 10^{-8}$. We note that H$_2$CO column densities of $\\sim 7\\times 10^{12}\\,$cm$^{-2}$ have been reported in diffuse clouds, including the ones discussed here (see Liszt et al. 2006).  Even though the value of $\\Delta L$ is not well constrained within a density range of $10^3$ to $10^5\\,$cm$^{-3}$, we can assess how reasonable the $\\Delta L = 1.0 \\times 10^4\\,$AU value is. As pointed out by Moore \\& Marscher (1995), the extended brightness distribution of 3C111 enables a study of spatial changes of the H$_2$CO absorption by comparing spectral lines toward the NE, Core, and SW continuum sources. Detection of absorption along the three sight-lines shows that the overall cloud is more extended than the projected separation of the NE and SW components, i.e., $7\\times 10^4\\,$AU (0.3$\\,$pc) assuming a distance of 350$\\,$pc (Moore \\& Marscher 1995; obtained from a photometric/extinction analysis by Ungerechts \\& Thaddeus 1987)\\footnote{For consistency with Moore \\& Marscher (1995) we use the same distance they assumed for the foreground cloud, however, the possible association with the California Molecular Cloud would imply a distance of $\\sim$450$\\,$pc (Harvey et al. 2013).}. A $\\Delta L = 1.0 \\times 10^4\\,$AU would imply that the depth of the cloud is somewhat less than the extent of the cloud in the plane of the sky, which is consistent with having slightly different line parameters toward the different continuum sources (Table~3)\\footnote{We note that velocity components at $-$1\\kms~and $\\-2$\\kms~were detected toward the Core and the NE regions, whereas only a single broad line was detected in our lower signal-to-noise spectrum of the SW 3C111 region. Thus, the significantly different H$_2$CO SW 3C111 profile could be caused in part due to blended lines in the low signal-to-noise spectrum.}.  Assuming the same values of kinetic temperature (30$\\,$K), density ($10^4\\,$cm$^{-3}$), H$_2$CO abundance ($5\\times10^{-9}$), and abundance/density domains as above, we obtain a total molecular column density and slab thickness equal to $N(H_2) = 3 \\times 10^{20}\\,$cm$^{-2}$ ([$3 \\times 10^{19}$, $5 \\times 10^{21}$] cm$^{-2}$) and $\\Delta L = 2 \\times 10^3\\,$AU ([$6 \\times 10^1$, $1 \\times 10^{5}$] AU) for B0212+735. In the case of NRAO$\\,$150 we obtain $N(H_2) = 6 \\times 10^{20}\\,$cm$^{-2}$ ([$6 \\times 10^{19}$, $1 \\times 10^{22}$] cm$^{-2}$) and $\\Delta L = 4 \\times 10^3\\,$AU ([$2 \\times 10^2$, $2 \\times 10^{5}$] AU), while for BL$\\,$Lac we obtain $N(H_2) = 2 \\times 10^{20}\\,$cm$^{-2}$ ([$2 \\times 10^{19}$, $2 \\times 10^{21}$] cm$^{-2}$) and $\\Delta L = 10^3\\,$AU ([$3 \\times 10^1$, $5 \\times 10^{4}$] AU).  \\newpage  \\subsection{Physical Scales Probed by the Monitoring Observations}  The most important aspect to relate absorption variability to cloud structure is to be able to convert time-lapses to physical displacement of the clouds perpendicular to our sight line, i.e., the proper motion due to the combined effect of Earth's revolution and the motion of the Solar System and the cloud around the Galaxy. Marscher et al. (1993) and Moore \\& Marscher (1995) estimated that the magnitude of transverse displacement for the foreground clouds toward NRAO$\\,$150, 3C111 and BL Lac is of the order of $\\sim 4\\,$AU per year. Specifically, Moore \\& Marscher (1995) report that after approximately 3.4 years of monitoring observations, the foreground clouds toward NRAO$\\,$150 and BL Lac would have had transverse displacements of 13.9$\\,$AU and 12.4$\\,$AU, respectively. These estimates depend on the precise knowledge of the transverse velocity of the foreground clouds with respect to Earth, which has not been directly measured. Thus, we recalculated the transverse displacement of the foreground clouds including uncertainty in the cloud's peculiar velocity around the Galaxy, in addition to the line-of-sight displacements due to Earth's revolution, LSR reference frame correction, peculiar velocity of the Solar System and estimated Galactic circular rotation of the foreground clouds. To calculate the Galactic circular rotation, we assumed the distances reported by Moore \\& Marscher (1995), i.e., 700$\\,$pc (kinematic distance, Marscher et al. 1993), 350$\\,$pc (estimated from a photometric/extinction analysis by Ungerechts \\& Thaddeus 1987), and 330$\\,$pc (estimated from the average vertical height of Galactic molecular clouds in the solar neighborhood and the galactic latitude of the foreground cloud, Lucas \\& Liszt 1993) for the NRAO$\\,$150, 3C111, and BL$\\,$Lac foreground clouds, respectively. We used the rotation curve of Clemens (1985) for consistency with Marscher et al. (1993). To estimate the range of possible transverse displacements, we assume a peculiar velocity of the foreground clouds of $\\pm 10$\\kms~in the three galactocentric cylindrical coordinates. In a time lapse of 748 days [the time lapse of the Marscher et al. (1993) observations], we obtain a maximum displacement of $\\sim$10$\\,$AU and a minimum of $\\sim 0\\,$AU depending on the peculiar velocity of the clouds toward NRAO$\\,$150, 3C111 and BL$\\,$Lac. This range encompasses the displacements quoted by Marscher et al. (1993), but shows that the peculiar velocity of the foreground clouds has a significant effect on the displacement, and hence, on our ability to ``translate'' time-lapses into transverse displacements.  The new 2009 observations reported here enable us to study variability in a time lapse of almost two-decades. From 1990 to 2009, the maximum estimated displacement is $\\la 100\\,$AU and as small as its parallax motion ($\\la 1\\,$AU; i.e., the case when the peculiar velocity of the cloud results in no proper motion) for NRAO$\\,$150, 3C111, and BL$\\,$Lac. An estimate of the transverse displacement for B0212+735 is uncertain because there is no distance measurement to the foreground cloud reported in the literature, and kinematic distances (e.g., Watson et al. 2003) or statistical distances (e.g., Magnani et al. 1985) are unreliable given the Galactic latitude and the small V$_{LSR}$ velocity of the absorption. However, a similar displacement range is expected for the foreground cloud in B0212+735 because the magnitude of its $V_{LSR}$ is within that of 3C111 and NRAO$\\,$150 (Table~3), and the magnitude of its galactic latitude ($|b| = 12^o$) is similar to that of BL$\\,$Lac ($|b| = 10.4^o$) and 3C111 ($|b| = 8.8^o$).   \\subsection{Variability}  Figures~3 and 5 bottom right panels show the difference between the 2009 spectra and the previous observations for all sources (i.e, ``red'' minus ``black'' spectra of the corresponding left panel), which will be called {\\it residual spectra} hereafter. As mentioned above, clear variability is seen toward NRAO$\\,150$ and 3C111 Core, however the variability behavior observed toward both sources is quite different.  In the case of 3C111 Core, we found no long-term variability trend. The 2009 spectrum is the same within the noise to the September 1991 and December 1990 spectra, and only one channel from the 1994 spectrum differs by $\\sim$3.5$\\sigma$ with respect to the 2009 spectrum. Moore \\& Marscher (1995) reported variability of the H$_2$CO line from the July 1991 and December 1992 epochs with respect to their May 1994 spectrum. We also find variability above $3\\sigma$ when comparing our 2009 spectrum with the July 1991 and December 1992 data (Figure~5). Comparing the March 2009 spectrum to the December 1992 spectrum (D92$-$M09 residual spectrum in Figure~5), we see a line that resembles an inverse-P Cygni profile, which may be caused by a small velocity difference ($\\sim 0.04$\\kms; see Table~2) between the two different sight-lines sampled on March 2009 and December 1992. Evidence for such velocity shift was pointed out by Moore \\& Marscher (1995) when comparing their 1992 and 1994 spectra. The velocity difference is less than the FWHM of the line, thus, velocity gradients within the cloud could be responsible for the observed variability. We note that the M94-M09 residual spectrum shows a similar profile than the D92-M09 residual spectrum, albeit with significantly less signal to noise. We do not detect evidence for such residual line profile in the other epochs, nor toward the other sources, and at this point it is unclear why only the M92-M09 residual spectrum shows such a pronounced profile. Follow up observations of 3C111 are needed to detect more occurrences of velocity shifts and detailed analysis of synthetic spectra from computer simulations of turbulent clouds is needed to estimate the expected occurrence rate of such events.  Figure~5 also shows that the peak of the July 1991 spectrum is less deep than the March 2009 spectrum by $\\Delta e^{-\\tau} \\sim 0.025$; Moore \\& Marscher (1995) also report less absorption in the July 1991 spectrum with respect to the May 1994 spectrum by $\\Delta e^{-\\tau} \\sim 0.02$ (see their Figure~5), however the line in our residual spectrum is narrower than that from Moore \\& Marscher (1995). The 3C111 residual spectra do not show evidence of hyperfine structure due to low signal-to-noise, which as mentioned in the introduction, would strengthen the case for variability.  The difference between the July 1991 and March 2009 3C111 spectra could be explained by a change in the line-of-sight overall characteristics of the foreground cloud. Recalling that the main absorption feature in 3C111 Core is consistent with $n(H_2) = 10^4\\,$cm$^{-3}$, $T_K = 30\\,$K, $5\\times10^{-9}$ H$_2$CO abundance, $\\Delta L = 10^4\\,$AU, the difference between the July 1991 and March 2009 could be caused by an $\\sim$8\\% variability in overall H$_2$CO abundance without changing any other parameter. Smaller abundance gradients would be needed to explain the variability if the depth ($\\Delta L$) of the H$_2$CO absorption region would also have changed. Significant H$_2$CO abundance gradients in molecular clouds are know to occur, e.g., Young et al. (2004) found clear evidence of H$_2$CO depletion in the interior of pre-protostellar cores.  In the case of NRAO$\\,$150 (Figure~3), our new 2009 spectrum confirms the variability trend reported by Moore \\& Marscher (1995), i.e., the main absorption feature has been monotonically increasing since 1990. The case of NRAO$\\,$150 shows the advantages of using H$_2$CO absorption as a probe for variability against point sources, in particular, the residual spectra show detection of the $F=1-0$ hyperfine component and no variability of the $-17.3$\\kms~line, which, as mentioned above, demonstrate that the residual lines are true variability and not artifacts. Figure~6 shows the variability as a function of time. The new data point shows a change in the slope of the variability, suggesting that the 1990's observations sampled a region with a greater column density/abundance gradient than the region sampled by the 2009 data.  Assuming $n(H_2) = 10^4\\,$cm$^{-3}$, $T_K = 30\\,$K, $5\\times10^{-9}$ H$_2$CO abundance, $\\Delta L = 4 \\times 10^3\\,$AU and $N(H_2CO) = 3\\times 10^{12}\\,$cm$^{-2}$, the variability observed between 1990 and 2009 can be explained by an increase of $1.3\\times 10^3\\,$AU in the effective thickness of the molecular cloud in the line of sight of the quasar ($\\Delta L$) without any change in molecular density, temperature or abundance. As mentioned above, the transverse displacement of the line of sight across the cloud between 1990 and 2009 is $\\la 100\\,$AU. Assuming a transverse displacement of 50$\\,$AU (half the maximum expected displacement), and that the cloud has a cylindrical filamentary structure (circular cross section) with the sight line to NRAO$\\,$150 crossing the filament perpendicular to its length (Figure~7), we can calculate the radius of curvature of the circular cross section. We obtain a filament radius of $R \\sim 2 \\times 10^4\\,$AU, and find that the line-of-sight would be crossing the filament very close to its edge (see Figure~7). We stress that the value of the radius is representative of the physical scales and not a precise measurement because of the large uncertainties involved in $\\Delta L$ and transverse displacement, and the filament is unlikely to have a perfectly circular cross-section. Nevertheless, the molecular observations by Pety et al. (2008) indeed show that the line of sight toward NRAO$\\,$150 is crossing a filamentary structure, which has projected width of $\\sim 40$\\arcsec~(see their Figure~4, $-12 < V < -9.5\\,$\\kms~ range, 5.8\\arcsec~resolution), which at a distance of 700$\\,$pc would correspond to a radius of $\\sim 1.4\\times10^4\\,$AU. This value is similar to the radius of $\\sim 2 \\times 10^4\\,$AU calculated above. Thus, without changing any other parameter of the overall physical conditions of the gas (e.g., temperature, abundance, density) the variability observed toward NRAO$\\,$150 could be explained by a simple geometrical effect caused by the line-of-sight crossing the edge of a molecular filament.  In order to explain the variability in H$_2$CO absorption profiles, Marscher et al. (1993) and Moore \\& Marscher (1995) argue for the presence of high density ($10^6\\,$cm$^{-3}$) molecular clumps at scales of $\\sim 10\\,$AU. In contrast, here we have found that minor chemical (abundance) gradients and/or geometrical effects can account for the observed variability without the need for high density clumps. Also, high density clumps are not necessarily present in the case of foreground clouds that show no variability, such as toward B0212+735 or the $-17.2$\\kms~component toward NRAO$\\,$150. Moreover, we note that at such high densities ($10^6\\,$cm$^{-3}$) the H$_2$CO would be detected in emission instead of absorption against the CMB, which has never been observed in diffuse clouds.\\footnote{Thermal emission of the 6$\\,$cm H$_2$CO line has only been seen toward the Orion BN/KL region (Zuckerman et al. 1975, see also Araya et al. 2006) and non-thermal emission has been detected only toward young massive stellar objects (e.g., Araya et al. 2008).} Hence, we concur with the conclusions Liszt \\& Lucas (2000), that presented independent arguments against explaining absorption variability as caused by AU-sized high density inclusions in interstellar clouds. In particular, Liszt \\& Lucas (2000) found that the high-density clump/inclusion assumption is inconsistent with the weak HCO$^+$ emission in diffuse clouds. As discussed in their paper, in order to explain the weak HCO$^+$ emission, a low ratio of total (high-density) clump to cloud area would be needed, but this would be inconsistent with the frequent optical depth variations observed in HCO$^+$ profiles."
    },
    "1402/1402.0964_arXiv.txt": {
        "abstract": "{ We use the clustering properties of Luminous Red Galaxies (LRGs) and the growth rate data provided by the various galaxy surveys in order to constrain the growth index ($\\gamma$) of the linear matter fluctuations. We perform a standard $\\chi^2$-minimization procedure between theoretical expectations and data, followed by a joint likelihood analysis and we find a value of $\\gamma=0.56\\pm 0.05$, perfectly consistent with the expectations of the $\\Lambda$CDM model, and $\\Omega_{m0} =0.29\\pm 0.01$, in very good agreement with the latest Planck results. Our analysis provides significantly more stringent growth index constraints with respect to previous studies, as indicated by the fact that the corresponding uncertainty is only $\\sim 0.09 \\gamma$. Finally, allowing $\\gamma$ to vary with redshift in two manners (Taylor expansion around $z=0$, and Taylor expansion around the scale factor), we find that the combined statistical analysis between our clustering and literature growth data alleviates the degeneracy and obtain more stringent constraints with respect to other recent studies. } ",
        "introduction": "The statistical analysis of various cosmological data (SNIa, Cosmic Microwave Background-CMB, Baryonic Acoustic Oscillations-BAOs, Hubble parameter measurements etc) strongly suggests that we live in a spatially flat universe that consists of $\\sim 4\\%$ baryonic matter, $\\sim 26\\%$ dark matter and $\\sim 70\\%$ some sort of dark energy (hereafter DE) which is necessary to explain the accelerated expansion of the universe (cf. \\cite{Hicken2009,Komatsu2011,Blake2011,Hinshaw2013,Farooq2013,Ade2013,Sperg2013} and references therein). Although there is a common agreement regarding the ingredients of the universe, there are different views concerning the possible physical mechanism which is responsible for the cosmic acceleration. Briefly, the general path that one can follow in order to mathematically treat the accelerated expansion of the universe is to see DE either as a new field in nature or as a modification of General Relativity (see for review \\cite{Copeland2006,Caldwell2009,Amendola2010}).  An interesting approach to discriminate between scalar field DE and modified gravity is to use the evolution of the linear growth of matter perturbations $\\delta_{m} (z)=\\delta \\rho_m / \\rho_m$ \\cite{Linder2004,Linder2007,Mar2014}. Specifically, a useful tool in this kind of studies is the so called growth rate of clustering, which is defined as $f(a)=\\frac{d{\\rm ln}D}{d{\\rm ln}a}\\simeq \\Omega_{m}^{\\gamma}(a)$, where $a(z)=(1+z)^{-1}$ is the scale factor of the universe, $\\Omega_m(a)$ is the dimensionless matter density parameter, $\\gamma$ is the growth index and $D(a)=\\delta_{m}(a)/\\delta_{m}(a=1)$ is the linear growth factor scaled to unity at the present time \\cite{Peebles1993,Wang1998}. The accurate determination of the growth index is considered one of the main goals of Observational Cosmology because it can be used in order to check the validity of General Relativity (GR) on cosmological scales. The basic ingredient in this approach comes from the fact that $\\gamma$ depends weakly on the dark energy equation of state (hereafter EoS) parameter w$(z)$ \\cite{Linder2007}, implying that one can split the background expansion history, $H(z)$, constrained by geometric probes (SNIa, BAO, CMB), from the dynamical perturbations growth history. %  Theoretically speaking, it has been shown that for those DE models which are within the framework of GR and have a constant EoS parameter, the growth index $\\gamma$ is well approximated by $\\gamma \\simeq \\frac{3({\\rm w}-1)}{6{\\rm w}-5}$ \\cite{Linder2007,Silveira1994,Wang1998,Nesseris2008}. In the case of the concordance $\\Lambda$CDM model (w$(z)=-1$) the above formula reduces to $\\gamma \\approx 6/11$. Considering the braneworld model of \\cite{DGP2000} we have $\\gamma \\approx 11/16$ (see \\cite{Linder2007,Wei2008,Gong2008,Fu2009}). Finally, for some $f(R)$ gravity models it has been found that $\\gamma \\simeq 0.415-0.21z$ for various parameter values (see \\cite{Gannouji2009,Tsujikawa2009}), while for the Finsler-Randers cosmology, Basilakos \\& Stavrinos \\cite{Basilakos2013} found $\\gamma \\approx 9/14$.  From the large scale structure point of view, the study of the distribution of matter on large scales using different extragalactic mass tracers (galaxies, AGNs, clusters of galaxies etc) provides important constraints on structure formation theories. In particular, since gravity reflects, via gravitational instability, on the nature of clustering \\cite{Peebles1993} it has been proposed to use the clustering/biasing properties of the mass tracers in constraining cosmological models (see \\cite{Matsubara2004,Basilakos2005,Basilakos2006, Krumpe2013}) as well as to test the validity of GR on extragalactic scales (\\cite{Basilakosetal2012} for a recent review see \\cite{Bean2013}).  Based on the above arguments, the aim of the current study is to place constraints on the $(\\Omega_{m}, \\gamma)$ parameter space using the measured two-point angular correlation function (hereafter ACF) of the LRGs, % with known redshift distribution. The basic idea is to compare the measured and the theoretically predicted ACF, which is based on the 3D power spectrum, the evolution of the relevant bias parameter and the Limber's inversion integral equation. Therefore, the predicted theoretical ACF depends also on the growth of linear matter perturbations via which the $(\\Omega_{m},\\gamma)$ pair, which can thus be constrained. The merit of utilizing ACF data for such a task is related to the fact that we do not need to consider a fiducial cosmological model in order to derive the ACF data, as well as on the fact that the ACF is unaffected by redshift-space distortions. In addition to ACF, we use the recent growth rate data as collected by Nesseris \\& Garcia-Bellido \\cite{NesGa12}, Hudson \\& Turnbull \\cite{Hud12}, Beutler et al. \\cite{Beutler} and Basilakos et al. \\cite{BasilakosNes2013} in order to put tight constraints on $(\\Omega_{m},\\gamma)$.   The structure of the paper is as follows. In section 2 we present the angular correlation function data, measured for LRGs and the growth data. In section 3 we discuss the theoretical angular correlation function model and basic ingredients in order to calculate it, such as the linear growth of matter perturbations, the evolution of the linear bias factor and the CDM power spectrum. The details of our methodology used to fit models to the data and our results are presented in section 4, while our main conclusions in section 5. ",
        "conclusions": "In the epoch of intense cosmological studies aimed at testing the validity of general relativity on extragalactic scales, it is very important to minimize the amount of priors needed to successfully complete such an effort. One such prior is the growth index and its measurement at the $\\sim 1\\%$ accuracy level has been proposed as a necessary step for checking possible departures from general relativity at cosmological scales \\cite{Bean2013}. Therefore, it is of central importance to have independent determinations of $\\gamma$, because this will help to control the systematic effects that possibly affect individual methods and tracers of the growth of matter perturbations.  In this study we use the basic large scale structure properties such as the clustering of the 2SLAQ Luminous Red Galaxies together with the growth rate of clustering provided by the various galaxy surveys in order to constrain the growth index of matter perturbations. The results of the two analyzes are used in a joint likelihood fitting procedure which helps to reduce the parameter uncertainties. The outcome constraints are: $(\\Omega_{m0},\\gamma)=(0.29\\pm 0.01,0.56\\pm 0.05)$, which are the strongest (to our knowledge) joint constraints appearing in the literature. Also, we check that our growth results are quite robust against the choice of the transfer function of the power spectrum and the Planck priors which are available in the literature \\cite{Ade2013,Sperg2013}.  Finally, considering a time varying growth index: $\\gamma(z)=\\gamma_{0}+\\gamma_{1}X(z)$, with $X(z)=z$ or $X(z)=z/(1+z)$ we find, as all similar studies, that $\\gamma_{1}$ and $\\gamma_{0}$ are somehow degenerate. However, based on the joint statistical analysis we have managed to put tighter constraints on $\\gamma_{0}$. Although, we have reduced significantly the $\\gamma_{1}$ uncertainty with respect to previous studies, the corresponding error bars remain quite large. Future, dynamical data are expected to improve even further the relevant constraints (especially on $\\gamma_{1}$) and thus the validity of GR on cosmological scales will be effectively tested.   \\vspace{0.0cm}"
    },
    "1402/1402.2010.txt": {
        "abstract": "{We present for the first time the evolution of the photospheric electric currents during an eruptive X-class flare, accurately predicted by the standard 3D flare model. We analyze this evolution for the February 15, 2011 flare using HMI/SDO magnetic observations and find that localized currents in \\J-shaped ribbons increase to double their pre-flare intensity. Our 3D flare model, developed with the OHM code, suggests that these current ribbons, which develop at the location of EUV brightenings seen with AIA imagery, are driven by the collapse of the flare's coronal current layer. These findings of increased currents restricted in localized ribbons are consistent with the overall free energy decrease during a flare, and the shape of these ribbons also give an indication on how much twisted the erupting flux rope is.} Finally, this study further enhances the close correspondence obtained between the theoretical predictions of the standard 3D model and flare observations indicating that the main key physical elements are incorporated in the model.  { %} % { } % { } % { } % { } % % % % % ",
        "introduction": "\\label{sect_intro}  Eruptive flares are amongst the main drivers of space weather as they inject solar material under the form of coronal mass ejections into the interplanetary medium \\citep[\\eg ][]{Bothmer1994,Forbes2006}. These blasts of solar plasma are regularly at the origin of magnetic storms on Earth and other planets \\citep[see][]{Gosling1991,Gonzales1991,Prange2004}. Solar flares are classified in function of the peak flux of X-rays measured near Earth, and divided into different categories. The most energetic events are classified as X-class flares and have intensities starting from 10$^{-4}$~W/m$^2$.  The majority of X-class flares have been shown to originate from eruptive flares \\citep{Wang2007}, making them important events to be understood.  Eruptive flares have several observational characteristics, including the flare loops \\citep[e.g.][]{Schmieder1995} and twisted magnetic field lines, identified as the flux rope at the origin of coronal mass ejections \\citep[e.g.][]{Cheng11, Zhang2012, Patsourakos2013}. Flare ribbons are also important flare observational characteristics, as they are the locations for the largest radiative emission increase of flares, in extreme ultraviolet (EUV) and in soft X-rays \\citep[SXR, see e.g.][]{Schmieder1987, delZanna2006}. They are {a consequence of} the impact at the chromosphere of energetic particles launched from the coronal reconnection site \\citep[e.g.][]{Reid2012}. During an eruptive flare, a two-ribbon structure is often observed \\citep[e.g.][]{Asai2003,Ding2003}, with a typical \\J-shape for both ribbons in the different magnetic polarities {\\citep[see e.g.][]{Moore95,Chandra09,Warren2011}.}  Although the CSHKP 2D model \\citep{Carmichael1964,Sturrock1966,Hirayama1974,Kopp1976} and its further extensions \\citep[\\eg ][]{Forbes1986,Shibata1995} explain some of the eruptive flare characteristics, it fails at rendering their intrinsic 3D nature, and therefore the complex evolution of {the flare loops, the} flux rope, and the \\J-shaped ribbons.  As such, different mechanisms for eruptive flares have been proposed \\citep[see][ and references therein]{Forbes2006,Amari2003a,Amari2003b,Shibata2011,Aulanier2014}. One of these scenarios involves a torus-unstable flux rope \\citep{Bateman78,Lin1998,Kliem2006}, and numerical simulations of such an evolution have been able to reproduce many of the eruptive flare characteristics \\citep[\\eg ][]{Aulanier2012,Janvier2013,Dudik2014}. The expansion of the unstable flux rope leads to a feedback process as a thin current layer, formed underneath the flux rope, feeds the growing structure with twisted field lines via magnetic reconnection. This process is also at the origin of flare loops and their strong-to-weak shear evolution \\citep[\\eg ][]{Warren2011}.  This was quantified during the 2011, May 9 eruptive flare and was properly predicted in the present standard 3D model \\citep{Aulanier2012}.  The strong distortions of the magnetic field line mapping, referred to as Quasi-Separatrix Layers \\citep[QSLs,][]{Demoulin1996a, Demoulin1997} are quantified by the squashing degree parameter $Q$ \\citep{Titov02,Pariat2012}.  Strong current density regions develop at QSL locations \\citep{Aulanier05b,Janvier2013}. The QSLs are therefore preferential locations for ideal MHD to break down, \\ie\\ for reconnection to occur. Indeed, both QSLs and current density structures show similar time evolution and shapes as flare ribbons (see Figures~3 and 15 in \\citealt{Dudik2014} for the QSLs, and the outward motion of current ribbons shown in Figure 6 of \\citealt{Aulanier2012}).  \\begin{figure} \\centering \\includegraphics[bb=10 05 720 400,width=0.49\\textwidth,clip]{fig_1_flareoverview} \\caption{ AIA image in the 335~\\AA\\ filter showing the South-West quarter of the Sun during the peak of the 2011, February 15 X-class flare. This eruptive flare occurred in the central part of AR 11158, {close to the center disk. The center of the signal saturation and the diffraction pattern at the flare maximum marks the very localized flare region in spite of its large magnitude. } } \\label{figAIAoverview} \\end{figure}   %Then, if the photosphere satisfies line-tied conditions, the flare ribbons would be created at the footprints of QSLs, and photospheric currents would appear at similar locations as flare ribbons. This is however a drastic condition, that can be set in a numerical environment but that does not reflect the reality of the solar photosphere.  As such, the present paper aims at quantifying the evolution of the photospheric electric currents and at directly comparing it with that predicted with the 3D numerical extensions to the standard flare model. More precisely, we study the vertical component of the electric current vector that we will simply denote as the vertical current in the following. %Note that these 3D numerical extensions do not involve chromospheric/photospheric fluctuations and the whole model being set with the plasma $\\beta$=0, we should expect rather large discrepancies between the model predictions and the observations. %An analysis of the photospheric {vertical} currents has been proposed by \\citet{Guo2013} for the electric currents of coronal jets, and proved to be successful in understanding the underlying reconnection processes feeding consecutive jet events. %In the present paper, we quantify the evolution of photospheric vertical currents to gain further insight on the reconnection process predicted by our 3D standard flare model. This is made possible by the high cadence and good spatial resolution of the HMI instrument onboard of the Solar Dynamics Observatory (SDO), and with the analysis of the vector magnetic fields at the photospheric level.  %Analysis of photospheric structures also requires an observation of an active region that is rather near the central region of the solar disk, so as to avoid the degradation of observations due to projection effects. The February 15, 2011 event, occurring in active region 11158, is a very good candidate, for the strength of the event (X2.2) and for its location near the center disk (Figure~\\ref{figAIAoverview}). It is also a good candidate because such a compact active region flare allows better measurements of the horizontal magnetic field vector in the related sunspots, the magnetic field being more intense in those regions.    The remainder of the paper is as follows. In Section~\\ref{Sect:2}, we present a general overview of the February 15 flare and the SDO instruments, HMI and AIA. We also introduce the {UNNOFIT} inversion method that allows the calculation of the three components of the magnetic field, and therefore the photospheric vertical currents. In Section~\\ref{Sect:3}, we study {quantitatively} the evolution of the photospheric current ribbons. These results are compared in Section~\\ref{Sect:4} with the evolution of flare ribbons seen with the AIA instruments, as well as with the  current ribbons predicted in the 3D standard flare model developed in a numerical simulation. Finally, the implications of the present results are discussed in Section~\\ref{Sect:5}, followed by a summary and a conclusion in Section~\\ref{Sect:6}. %________________________________________________________________ ",
        "conclusions": "\\label{Sect:6} % % The present paper quantifies for the first time the photospheric current evolution during an eruptive flare and challenges the general understanding of the current density decrease associated with the evolution of the flare. Indeed, as the flare occurs, the magnetic field configuration becomes more potential, so that associated currents should globally decrease (see the discussion in Section \\ref{Sect:5}). The 3D extensions of the standard eruptive flare model has successfully predicted some observational characteristics of eruptive flares, as presented in a series of papers {(see Section~\\ref{sect_intro}).} Especially, the model shows that current ribbons form at the photospheric boundaries, evolve with similar patterns to the flare ribbons, and are associated with the sudden collapse of a coronal current layer which further thinning is driven by the torus-unstable flux rope.  In this work, we aim at comparing the evolution of currents during an eruptive flare so as to confront with the results predicted with the 3D standard flare model. To do so, we have analyzed the vertical current density obtained from HMI magnetograms and with the UNNOFIT inversion code during the X-class flare event of February 15, 2011. The other components of the current density cannot be computed due to the single spectral line of the HMI observations. However, the vertical component of the current density is of highest interest, since it is used in magnetic extrapolation of coronal magnetic field, and since this component can be injected in the corona. {Also, the 3D standard model indicates that the regions of the photospheric current ribbons are nearly force-free. The magnetic field vertical component being dominant, it results that the current density vector $\\bf{J}$ is nearly $\\bf{B}$-aligned, and therefore dominantly vertical too.} The vertical current density is dominated by direct currents, especially in the flare region. The current maps and the base-difference of the direct currents are used to follow the evolution of the current ribbons.  {The paper then reports on the following main findings for the observation of the evolution of the vertical current density:}  \\begin{itemize} \\item{{During the impulsive phase, the current density increases in the hooks as well as in the straight parts of the \\J-shaped current ribbons. The spatial and temporal coherence of the current ribbons shows that they are unlikely to be related to polarization artifacts. }  } \\item{{The quantification of the electric current shows increases by about 1.5 to 2 times that remains more than one hour after the occurrence of the flare peak.} } \\item{{The structural changes of the current ribbons as seen in the observations are the same as found in the photospheric currents of the 3D MHD simulation. The two main characteristics are: first the expansion of the hook ribbon regions, that become rounder as the flux rope continues to grow, and second, the elongation of the straight part of the ribbons. } } \\item{{The comparison of the vertical current component, in \\Jz-maps, with the flare ribbons obtained from AIA 335~\\AA\\ filter shows similar evolutions of the current and the flare ribbons, demonstrating that {both current and flare} ribbons result from the evolution of the magnetic field.} } \\end{itemize}      These results therefore confirm the predictions of the 3D MHD simulation of an eruptive flare and that despite not being a data driven simulation, it still incorporates at the lower boundary the key typical ingredients of an active region evolution with theoretical boundary conditions. In particular, it shows that the approximation of line-tied magnetic field at the photospheric boundaries is correct, as the horizontal magnetic field is still allowed to evolve.  The model allows a further understanding on the evolution of the photospheric current. A current layer forms all along the QSLs, from the corona to the photospheric layer, where the footprints can be traced as the current ribbons.  {The ideal instability of the flux rope, with the photospheric line-tying, further enhances the electric current within the current layers as it imposes a different evolution of the magnetic field on both side of QSLs.  Such fast} evolution ultimately leads to the collapse of the current layer and a fast increase of the reconnection rate.  This explains the impulsive phase of flares.  The results also have implications on the magnetic configuration involved during an eruptive flare, as follows. The shape of flare/current ribbons is consistent with the presence of a flux rope with a twist of less, or equal to one turn. Indeed, since these ribbons follow the shape of the QSLs (see cartoon in Figure \\ref{figmodel}), a more twisted flux rope structure would lead to a more complicated whirl shape for the photospheric trace of QSLs (see for example Figure 8 in \\citealt{Demoulin1996b}), and therefore not to a \\J-shaped structure as seen here for current ribbons ($J_z$-maps) and flare ribbons (335 \\AA ).  {The observed ribbon shapes therefore exclude the possibility to have more complicated, highly twisted, kink unstable flux ropes, as is required in some models of CME ejection (\\textit{e.g.} \\citealt{Gibson1998}; \\citealt{Low2001} (Fig.9); \\citealt{Roussev2003, Cohen2010, Kliem2010, Shiota2010, Lugaz2011, Roussev2012}). It should be noted that, however, those studies focus on the propagation of CMEs and not on the flare. Nevertheless, it is arguable that the development of a consistent eruption model, with the flare and the CME up to its interplanetary propagation, need to match with observed \\J-shaped current and flare ribbons, since they indicate how twisted the CME magnetic field is. This is arguably important, for example because the CME acceleration depends on the Lorentz force, which itself depends on the currents, and also because the reconnection of the twisted fields with the ambient fields during the early CME expansion strongly depends on their relative orientations. }  The fact that the current ribbons remain coherent at the photosphere demonstrates that the chromosphere does not have a strong influence on the paths of the currents from the coronal reconnection region towards the photosphere. This was not straightforward a priori. Indeed, the increasing plasma $\\beta$, through the dynamic chromosphere down to the thick photosphere, must lead to the amplification of non-negligible cross-field currents. Nevertheless, the HMI observations show that the $\\beta$=0 and the infinitely thin photosphere hypotheses in the present MHD numerical simulation remain a good approximation.  {In conclusion, the electric current evolution, in the studied flare of February 15, 2011, is in agreement with the predictions made by the 3D standard model developed throughout a series of papers with the OHM code. The same model explains the evolution of flare loops \\citep{Aulanier2012}, the underlying 3D reconnection process \\citep{Janvier2013}, and the slipping motion of field lines during an eruptive flare event \\citep{Dudik2014}. This model also quantifies the energy release in flares in function of the spatial size and magnetic flux of active regions \\citep{Aulanier2013}.  Finally, all these correspondences between this 3D standard model and observations are strong evidences that the key physical ingredients are included in the model, and confirm the hypothesis of line-tied, $\\beta$=0 conditions given to the model.   % % %"
    },
    "1402/1402.2975_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec_intro} Understanding the formation and evolution of the first generation of stars would help in fathoming the initial chemical conditions and evolution of the universe \\citep{brom04}. In addition, it is believed that the radiation from such Population III stars may have played an important role in the reionization of the universe. These massive stars emit 10$^{47}$ $-$ 10$^{48}$ $\\hon$ and $\\heo$ ionizing photons per second per solar mass of formed stars, where the lower value holds for stars of $\\sim$ 20 M$\\odot$ and the upper value applies to stars with $\\gtrsim$ 100 M$\\odot$ \\citep[see][]{tum00,brom01}. However, their lifetimes are short ($\\sim$ 3 $\\times$ 10$^{6}$ years) and hence it would be very difficult to observe them directly. Nevertheless, we can try to infer signatures of Population III stars in the most metal-poor environments detected in nature. Recently there have also been indications of Population III star formation in the universe, even as late as z $\\sim$ 3, possibly due to inefficient transport of heavy elements and/or poor mixing that leaves pockets of pristine gas \\citep{jim06,tor07,ino11,cas13}. \\\\ \\indent The metal-poor stars in the halo of the Milky Way are being studied in an effort to understand the chemical composition and hence the nature of star formation at very early epochs \\citep[see][]{akr04,cay04,beer05,freb10}. Such studies can also be used to infer the chemical evolution in the universe, if most galaxies follow the chemical history of the Milky Way. However, the measurement of abundances in stellar atmospheres is hampered by having to take into account non-local thermodynamic equilibrium effects as well as three dimensional effects in one dimensional stellar atmosphere models \\citep{asp05}. In addition, the presence of convection or accretion of material from a companion can affect the photospheric abundance measurements. Hence, these measurements may not always trace well the chemical evolution history of the Galaxy. \\\\ \\indent These issues are not present when one studies damped Ly$\\alpha$ systems (DLAs), which are by definition, clouds of neutral gas with neutral hydrogen column density \\textit{N}($\\hon$) \\textgreater~2 $\\times$ 10$^{20}$ $\\hon$ atoms cm$^{-2}$ \\citep[see for a review] []{wolf05}. The high \\textit{N}($\\hon$) self-shields the gas from the ultraviolet background radiation of quasars (QSOs) and galaxies. Due to this, abundance measurements require negligible ionization corrections, i.e., hydrogen gas in DLAs is mostly neutral and most of the heavy elements are either neutral or singly ionized \\citep{petn02,pro02,cen03,led03}. Even then, line saturation and presence of dust may hinder abundance measurements in DLAs. However, when the metallicity of the DLAs happens to be $\\lesssim$ 10$^{-2}$ Z$_{\\odot}$ (i.e. in so called metal-poor DLAs), these issues are also mitigated \\citep[for e.g.][]{akr05}, especially for the estimation of C, N and O abundances. Moreover, high-z metal-poor DLAs are believed to probe gas in or around protogalaxies, and hence contain evidences of the earliest star formation and retain signatures of the initial chemical enrichment. The relative abundances of C, N, O, Fe, etc., in low-metallicity DLAs can give insights into the type of stars (specifically the stellar initial mass function, IMF) that led to their production, and hence shed light on our understanding of galactic nucleosynthesis and chemical evolution \\citep{petj08,petn08, penp10,cookb11}. In view of this, we undertake a study of five DLAs (pre-selected as candidate low-metallicity DLAs), and their elemental abundances, and compare our results with the existing measurements of low-metallicity DLAs and stars in the Galactic halo.\\\\ \\indent There is an ongoing interest in detecting and studying the properties of metal-poor DLAs. \\citet{petn08} studied a sample of four metal-poor DLAs (defined as those with [Fe/H] $\\le$ $-$2.0) using high resolution spectra (R $\\sim$ 40000) from Keck and VLT. They found that the C/O ratio trend at low metallicity in high redshift DLAs matches that of halo stars in the Galaxy \\citep{akr04,spi05}, and the ratio shows higher values that cannot be explained by nucleosynthesis models based on Population II stars. Their results also suggested that the N/O ratio at low metallicities may show a minimum value. \\citet{penp10} used medium resolution spectra (R $\\sim$ 5000) to study a sample of 35 DLAs pre-selected as metal-poor from the SDSS DR5 database. However, while medium resolution spectra is sufficient to detect metal-poor DLAs, one needs to be cautious about employing them to measure accurate abundances in the most metal-poor systems, which typically have line widths $\\lesssim$ 10 km s$^{-1}$. \\\\ \\indent Measurements of seven new metal-poor DLAs observed with high resolution spectrograph were reported by \\citet{cookb11}, who have also compiled a sample of all the (22) metal-poor DLAs known in literature. Using this sample and taking oxygen as the reference element, they have defined the typical abundance pattern of a very metal poor (VMP) DLA by determining the mean $\\langle$X/O$\\rangle$ ratio for each available element X, and then referring this mean value to the adopted solar scale. On comparing this with the yields of both Population II and III stars, it is found that, while a standard model of Population III stars (a top-heavy IMF where the stars explode as core collapse supernovae) gives reasonable agreement with the observed abundance pattern of a typical VMP DLA, the possibility that Population II stars may account for the observed abundances cannot be discarded. To distinguish between these two nucleosynthesis models, the [C/Fe] ratio is useful, as current models of Population II stars cannot explain an enhanced [C/Fe] ratio (along with [N/Fe]$\\lesssim$ 0). An enhanced [O/Fe] ratio ($\\gtrsim$ 0.4) will also be an useful discriminatory tool, since enhanced [C/Fe] is likely to lead to enhanced [O/Fe]. There are only two examples of carbon enhancement in metal-poor DLAs \\citep{cooka11,cook12}, which are likely to be the outcome of nucleosynthesis in massive stars. Understanding the abundance pattern in these carbon-enhanced metal-poor (CEMP) DLAs and increasing their number is important for understanding the true metal yields from very massive stars. \\\\ \\indent In addition, since DLAs are the main reservoirs of neutral gas at high z \\citep{not09a,pro05}, studying their star formation properties is important, as they could contribute significantly to the global star formation rate density at high z \\citep[for e.g.,][] {wolf08,rah10}. Detection of H$_{2}$ or $\\cto$* absorption in DLAs can allow us to study the physical conditions in the absorbing gas, and also in principle to infer the star formation rate \\citep[see][]{wolf03a,sri05}. In the present sample, we have observed $\\cto$* absorption associated with two of the DLAs. The $\\cto$* $\\lambda$1335 column density has been used to measure the [$\\cto$] $\\lambda$158 $\\mu$m cooling rate in the neutral gas. We simulate the physical conditions in these two DLAs, and try to match the model predictions with the observations in order to infer the ambient radiation field and cosmic ray ionization rate in these DLAs. Further, we look at all the detections of $\\cto$* in metal-poor DLAs reported till date and try to relate the $\\cto$* absorption to the metallicity and \\textit{N}($\\hon$), in order to understand the physical conditions in such DLAs. \\\\ \\indent This paper is organized as follows. In Section 2, we give details of the observations and data reduction process. How the analysis of the data was carried out is explained is Section 3. Section 4 gives description of the individual absorption systems. In Section 5, we discuss the trends observed in abundance ratios of elements in metal-poor DLAs. We discuss the detection of $\\cto$* in metal-poor DLAs and its implications in Section 6. Lastly, we summarise our results and present the conclusions in Section 7. ",
        "conclusions": "\\label{con} We have studied a sample of five DLAs and one sub-DLA and obtained their elemental abundances. Four of these DLAs and the sub-DLA turn out to be metal-poor, with [Fe/H] $\\lesssim$ $-$2.0, and among these DLAs, two are extremely metal-poor, with [Fe/H] $\\simeq$ $-$3.0. One of the extremely metal-poor systems (z$_{abs}$ = 2.34006 towards J0035$-$0918) studied here has also been analysed by \\citet{cooka11}. This is also one of the two instances of CEMP DLAs that have been detected. However, we find that the extremely large enhancement of C in this system, as reported by \\citet{cooka11}, is most probably due to the $b$ parameter being constrained to be the same for both the high and low mass ions, or in other words neglecting the thermal broadening. From our curve of growth analysis and best fit to the metal line profiles using VPFIT, we come to the conclusion that the line widths are likely to be determined by thermal broadening. In this case, carbon is still enhanced over iron, but by about three times. We estimate the DLA to have temperature in the range of 5000 $-$ 8000 K. Using this and from our CLOUDY model, we find that the physical conditions in this gas are consistent with a WNM in ionization equilibrium with the metagalactic UV radiation field dominated by the QSOs and galaxies. \\\\ \\indent The system at z$_{abs}$ = 4.20287 towards J0953$-$0504 is the most metal-poor DLA at z $>$ 4 studied till date. It shows no enhancement of carbon or oxygen and its relative elemental abundance pattern is similar to that of a typical metal-poor DLA, which may be explained by standard Population II nucleosynthesis results. The O/Fe, N/O and C/O ratios measured for our full sample conform to the observed trends seen in other metal-poor DLAs and stars. \\\\ \\indent In addition, we have studied the cooling and heating processes in two of the DLAs showing $\\cto$* absorption using CLOUDY. The $\\cto$* absorption is found to arise in the CNM phase and accounts for $\\gtrsim$ 80\\% of the total cooling in the DLAs. For the CNM solution, heating by cosmic rays is found to contribute $\\sim$ 60\\% to the total heating in the DLAs. Further, we find that an intergalactic background radiation field is not sufficient to produce the observed $\\cto$* cooling in the DLAs, implying in-situ star formation in the DLAs, which is the same conclusion as reached by \\citet{wolf06} and \\citet{wolf08}. We infer the radiation field (in particular the CR ionization rate) in the DLAs to be similar to the local interstellar medium with amplitude scaled by some factor (G$_{0}$ ranging from 1 to 4 times that of the Galaxy). \\\\ \\indent Lastly, we have assembled a sample of DLAs with [Fe/H] $\\le$ $-$2.0 and showing $\\cto$* absorption from the literature. Using photo-ionization models we argue that in such low-metal and low-dust systems, heating by grains is not as effective as in the Galactic ISM, and cosmic rays are most probably responsible for the observed excitation of $\\cto$*. The inferred CR ionization rate varies from same as (in case of ``low cool'' metal poor DLAs) to more than four times (in case of ``high cool'' metal poor DLAs) that seen in the Milky Way. Understanding this high CR-ionizing rate in the metal poor DLAs will enable us to understand star formation history, role of infall/winds and magnetic fields in low-metallicity galaxies in the early universe.  \\noindent"
    },
    "1402/1402.5880_arXiv.txt": {
        "abstract": "\\noindent We develop a non-singular bouncing cosmology using a non-trivial coupling of general relativity to fermionic fields. The usual Big Bang singularity is avoided thanks to a negative energy density contribution from the fermions. Our theory is ghost-free since the fermionic operator that generates the bounce is equivalent to torsion, which has no kinetic terms. The physical system consists of standard general relativity plus a topological sector for gravity, and fermionic matter described by Dirac fields with a non-minimal coupling. We show that a scale invariant power-spectrum generated in the contracting phase can be recovered by suitable choices of fermion number density and bare mass, thus providing a possible alternative to the inflationary scenario. ",
        "introduction": "\\noindent The pioneering work of Hawking and Penrose demonstrated that at initial times, the Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) metric of the Standard Big Bang cosmology suffers from singularities in all curvature invariants\\cite{HE}. Their theorem states that the initial singularity is unavoidable if space-time is described by General Relativity and if matter obeys the null energy condition. Over the years non-singular bouncing cosmologies have been proposed to avoid the Big Bang singularity by obviating one or all of the assumptions behind the Hawking-Penrose theorem.  However, a successful theory of the early universe must predict the observed nearly scale-invariant spectrum of adiabatic fluctuations in the CMBR.  Scale invariance was attempted in the context of bouncing models with a contracting phase such as Ekpyrotic \\cite{Ekpy}, String Gas \\cite{stringas} and Pre-Big-Bang scenarios \\cite{BF}.  On the other hand, it has proven difficult to obtain adiabatic scale invariant fluctuations in the contracting phase in a number of these models, mainly due to issues in mode matching between the contracting and expanding phases \\cite{BF}.  In pioneering works by Brandenberger, Finelli and independently by Wands \\cite{BF,Wands}, it was shown that it is possible to generate a scale-invariant power spectrum in a matter dominated contracting universe.  These authors demonstrated a ``duality'' between the scale invariant power spectrum generated in the inflationary epoch and a contracting matter dominated phase.  During the contracting phase, gauge-invariant perturbations that cross the Hubble-scale are scale-invariant if the scale factor evolves as $a(t) \\!\\sim\\! (-t)^{2/3}$.  Furthermore, if the bounce is non-singular, the scale-invariant modes can be matched to scale-invariant modes in the expanding phase.  A handful of matter bounce scenarios has since been proposed, mostly based on fundamental scalar fields \\cite{quintom, ghost1, ghost2, Ekpy, ek2}. In this paper we present a matter bounce scenario based on Dirac fermions, specifically, on the four-fermion interaction. This line of research has been previously addressed in Ref.~\\cite{AP} from the perspective of a wide class of generic potentials of the Dirac field.  Our work differs from Ref.~\\cite{AP}, where the four-fermion interaction was neglected in computing the scalar perturbation and a de Sitter background expansion was assumed.  In this work we follow the perspective taken by ~\\cite{AB, Pop} where a torsion induced four-fermion interaction yields a non-singular bounce. In Ref.~\\cite{Magueijo:2012ug}, the role of a parity violating 4-fermion self-interaction in a torsion-free theory has been studied. In this work, we propose both a non-minimal coupling of fermions (as analyzed in Refs~\\cite{M, BD, BD2}) and a topological gravitational term endowed with torsion (see for instance Ref.~\\cite{PR}). We show that these two terms generate four-fermion interactions whose dynamics yield a scale-invariant power spectrum\\footnote{The presence of the torsion background turns the topological term in the Holst action from a surface term into a contribution to the four-fermion interaction term.}.  Furthermore, the four-fermion current density modifies the Friedman equations to have a negative energy density that redshifts like $\\sim a(t)^{6}$. We show that the resulting bounce is non-singular provided that anisotropic stress is sub-dominant\\footnote{In a companion paper \\cite{paper00}, some of us used these findings as a starting point to discuss consequences for the fate of black hole solutions \\cite{collasso}, which for a  suitable choice of some parameters of the theory may never form.}, consistent with previous literature. Moreover, we show for the first time that the adiabatic quantum fluctuation of fermions in the contracting phase can be scale-invariant. Indeed, as we may straightforwardly infer from the result of Brandenberger and Finelli, since the bounce is non-singular our scale-invariant curvature perturbation induced by the fermion quantum fluctuations will enter the expanding phase as a scale invariant fluctuation.  An advantage of the mechanism shown in this paper is that it does not require any fundamental scalar field. The fermionic field is sufficient to account for both the bounce and the generation of nearly scale-invariant scalar perturbations.  The following is an outline of the paper. In Section II we introduce the theoretical framework, and cite the relevant works in the literature. In Section III we address the consequences of the model for the Matter-Bounce scenario. In Section IV we address the cosmological perturbations induced by the fermionic field. In Section V we discuss consistency with experimental data. In Section VI we provide some concluding remarks and mention works in progress. ",
        "conclusions": "\\noindent When general covariance accommodates non-minimal coupling in the fermionic sector, a four-fermion interaction modifies the cosmological evolution to yield a bounce.  In this work we have demonstrated that the same fermions that regulate the singularity also generate scale-invariant quantum fluctuations in the contracting phase. Using the arguments of Brandenberger and Finelli, we can easily match these fermionic perturbations to the scale invariant modes in the expanding phase.  The bounce is non-singular because the torsion, which is responsible for the bounce, does not receive quantum corrections. The gravitational wave power spectrum and the resulting tensor to scalar ratio have been derived. In a future paper, it will be interesting to compute corrections to the tensor to scalar ratio due to the coupling of the gravitons to the fermions.  Furthermore, in order to fully understand the generation of scale-invariant scalar perturbations, it is essential to address the mechanism here discussed in terms of the canonical Mukhanov-Sasaki variables, to which both matter and metric perturbations contribute. In a forthcoming paper \\cite{nuova} some of us are considering how to recover those variables by looking at the second order action of the theory.  In this context, it is interesting to notice that while matter perturbations are derived from the relevant fermionic bilinears, which behave as scalar and vector fields, perturbations of the fermionic bilinears must be expressed in terms of the fundamental fermionic fields, the dynamics of which is dictated by first order differential equations. This feature is at the origin of the very different behavior of fermionic matter perturbations relative to scalar field perturbations."
    },
    "1402/1402.5414_arXiv.txt": {
        "abstract": "Dynamically new comet C/2011 L4 (PanSTARRS) is one of the brightest comets since the great comet C/1995 O1 (Hale-Bopp). Here, we present our multi-wavelength observations of C/2011 L4 during its in-bound passage to the inner Solar system. A strong absorption band of water ice at 2.0 $\\mu$m was detected in the near infrared spectra, taken with the 8-m Gemini-North and 3-m IRTF telescopes. The companion 1.5~$\\mu$m band of water ice, however, was not observed. Spectral modeling show that the absence of the 1.5~$\\mu$m feature can be explained by the presence of sub-micron-sized fine ice grains. No gas lines (i.e.\\ CN, HCN or CO) were observed pre-perihelion either in the optical or in the sub-millimeter. 3-$\\sigma$ upper limits to the CN and CO  production rates were derived. The comet exhibited  a very strong continuum in the optical and its slope seemed to become redder as the comet approached the Sun. Our observations suggest that C/2011 L4 is an unusually dust-rich comet with a dust-to-gas mass ratio $>$ 4. ",
        "introduction": "Comet C/2011 L4 (PanSTARRS) was discovered at 7.9 AU from the Sun in June 2011 by the 1.8-m Pan-STARRS 1 survey telescope atop Haleakala. The detection was confirmed in followup observations with the Canada-France-Hawaii Telescope \\citep{wainscoat:2011}. The small reciprocal of the original semi-major axis, \\textit{1/a}= -8.9$\\times$ 10$^{-5}$ AU$^{-1}$\\citep{williams:2013}, suggesting that this comet is a recent arrival from the Oort Cloud. C/2011 L4 reached an apparent visual magnitude of \\textrm{-1} at 0.301 AU in March 2013, which made it one of the brightest comets in the past two decades since comet C/1995 O1 (Hale-Bopp). The apparition of C/2011 L4 provided a rare opportunity to monitor a dynamically new comet over a significant time period and a wide range of heliocentric distances, particularly during the in-bound leg.  At the time of discovery, C/2011 L4 showed an extended appearance \\citep{wainscoat:2011}, indicating the presence of a substantial coma. The turn-on distance for any appreciable water ice sublimation is around 5-6 AU from the Sun, beyond which the equilibrium surface temperatures is too low and H$_2$O-ice sublimation is not sufficient to lift even the smallest sub-micron-sized dust particles from the surface \\citep{meech:2004}. It is clear that C/2011 L4 was far beyond the water sublimating zone yet it appeared active. Several possible mechanisms that could power distant comet activity have been proposed, including latent heat release from the amorphous-to-crystalline water ice transition \\citep{prialnik:1992, notesco:2003, bar-Nun:2003}, sublimation of frozen super-volatiles \\citep{meech:2004} and comet fragmentation \\citep{boehnhardt:2004}. Among these, the exothermic amorphous-to-crystalline phase transition of water ice is considered to be an important energy source that can trigger distant cometary activity and power a cometary outburst \\citep{enzian:1997}. Observationally, crystalline ice differs from its amorphous counterpart by the presence of a sharp narrow feature at 1.65~$\\mu$m. This feature has been observed in large Kuiper belt objects via near infrared (NIR) spectroscopy \\citep{Jewitt:2004a, trujillo:2007} and was recently detected in the coma of an outbursting Jupiter-family comet (JFC) P/2010 H2 (Vales) \\citep{yang:2010}. However, crystalline water ice has never been seen in any Oort Cloud comet (OCC) \\citep{davies:1997, kawakita:2004}. Nevertheless, the absence of observational evidences does not exclude the presence of crystalline ice in OCCs. This is because the available OCC spectra have limited signal-to-noise ratios (SNR) and spectral resolutions. We conducted a multi-wavelength observational campaign of C/2011 L4, from optical to sub-millimeter wavelengths. The main goals of this study were three-fold: (i) search for water ice, in particular, the crystalline feature at 1.65~$\\mu$m via high SNR and medium-resolution spectroscopy in the NIR; (ii) monitor the on-set of volatile species (such as CN, HCN and CO, etc) in the optical and the radio as the comet approached the Sun, (iii) monitor the development of the dust and gas comae via optical spectroscopy. ",
        "conclusions": "\\subsection{Water Ice Features} The critical temperature for amorphous-to-crystalline ice transition is T$_c$ $\\sim$ 140 K, at which the timescale of crystallization is about 1 hour. The highest possible surface temperature for C/2011L4 should be that at its subsolar point, where T$_{ss}$(r$_h$)=T$_0$/$\\sqrt{r_h}$. If we assume C/2011 L4 is a slow rotator, then T$_0$=392 K \\citep{jewitt:2009}. At r$_h$=5.2 AU, C/2011 L4 can reach T$_{ss}\\sim$170K by solar heating alone. Although this is an extreme case, it shows that it is possible that at least part of the surface or near-surface ice could have experienced crystallization within our observation window.  As shown in Figure 2, sub-micron-sized ice grains show a very weak or non-existent absorbing feature at 1.65~$\\mu m$ regardless of the physical structure of individual grains. Thus, the absence of the 1.65~$\\mu m$ band does not exclude the presence of very fine crystalline ice grains. When focusing on the 2.0~$\\mu$m band, our $\\chi^2$-fit show that amorphous ice fits the comet spectra better than crystalline ice for the March and June data, while the crystalline model fits the July data marginally better. However, the small difference between these two ice models is within the noise level of our data. Therefore, we cannot constrain the crystallinity of the icy grains of C/2011 L4, mainly due to their small sizes.  Although the observed asymmetry in absorption strength between the bands at 1.5~$\\mu m$ and 2.0~$\\mu m $ can be successfully explained by sub-micron pure water ice grains, there are other explanations for the observed features. For example, the 2.0 $\\mu m$ absorption feature may not be solely generated by pure water ice grains but in addition be due to the superposition another absorption band of organic materials or hydrated minerals. We explored this possibility by searching for matching materials that exhibit a prominent absorption feature at or near 2.0~$\\mu$m. We investigated several commonly used spectral libraries, such as RELAB \\citep{hiroi:1996}, USGS \\citep{clark:2003} and ASTER \\citep{baldridge:2009}, which include examined materials from terrestrial minerals to meteorite samples. We find no other material with a NIR signature that matches the shape and position of the 2.0 $\\mu m$ band of C/2011 L4.  \\subsection{Dust Coma} Shown in Fig 3b, the slope of the dust continuum became steeper as the comet approached the Sun. Similar color-$r_h$ trends have been reported in other comets \\citep{hartmann:1984}. However, later studies \\citep{jewitt:1988} did not confirm these reports, suggesting instead that the reported trend may have been due to the contamination of cometary emissions. Given that no cometary emissions were observed in our optical data, the observed color trend is not caused by gas contamination. Possible explanations for this correlation, if real, include: (a) increase in the effective size of coma grains, (b) decrease in the ice-to-dust ratio in coma grains and (c) changes in phase angle as the comet approaches the Sun. Our NIR observations suggest that the ice-to-dust ratio remained more or less the same (within uncertainties) and the phase angle of C/2011 L4 did not change much from April to July. Therefore both (b) and (c) scenarios are unlikely and we are left with (a) as the only plausible mechanism: the observed color-$r_h$ trend is observed because larger dust particles being released into the coma as the comet moved closer to the Sun.  Given that no gas was detected, what could be the driving force for the cometary activity? One possibility is that C/2011 L4 underwent a strong outburst which released most of the observed dust and ice grains into the coma. At the time of our observations, the major gas production had ceased and the nucleus was surrounded by a remnant dust cloud. However, the change in the slope of the dust continuum argues for an on-going mass loss process. C/2011 L4 is a dynamically new comets, so reactive species could have been synthesized via long term exposure to cosmic rays or galactic UV bombardment while the comet was in the Oort Cloud \\citep{hughes:1990}. These radical species could explode when exposed to the solar wind and that could in turn trigger several small outbursts \\citep{niedner:1980}. Alternatively, the distant activity of C/2011 L4 could be driven by highly volatile ices, such as CO$_2$. Recently, several comets, such as comet 103P/Hartley 2 and comet ISON, haven observed to show strong CO$_2$ signatures \\citep{ahearn:2011, lisse:2013}. Unfortunately, CO$_2$ can not be observed from the ground and no space observations are available for C/2011 L4. To verify whether CO$_2$ could be driving the cometary activity, evolutionary models \\citep{meech:2013} which attempt to match the brightness measurements of C/2011 L4 are needed.  In summary, our multiwavelength observations yield the following findings: 1) very fine water ice grains in the coma of C/2011 L4; 2) no gas emission was detected; 3) C/2011 L4 is unusually dust-rich; 4) the dust coma seemed to becoming redder with decreasing r$_h$."
    },
    "1402/1402.1237_arXiv.txt": {
        "abstract": "Our understanding of massive star evolution is in flux, due to recent upheavals in our view of mass loss, and observations of a high binary fraction among O-type stars.  Mass-loss rates for standard metallicity-dependent line-driven winds of hot stars are now thought to be lower by a factor of 2-3 compared to rates adopted in modern stellar evolution models, due to the influence of clumping on observed diagnostics.  Weaker line-driven winds shift the burden of H-envelope removal elsewhere, so that the dominant modes of mass loss are the winds, pulsations, and eruptions of evolved supergiants, as well as binary mass transfer.  Studies of stripped-envelope supernovae, in particular, require binary mass transfer.  Dramatic examples of eruptive mass loss are seen in Type IIn supernovae, which have massive shells ejected just a few years before core collapse. These are a prelude to core collapse, and may signify severe instabilities in the latest nuclear burning phases. The shifting emphasis from steady winds to episodic mass loss is a major change for low-metallicity regions, since eruptions and binary mass transfer are less sensitive to metallicity.  We encounter the predicament that the most important modes of mass loss are also the most uncertain, undermining the predictive power of single-star evolution models beyond core H burning.  Moreover, the influence of winds and rotation in models has been evaluated by testing single-star models against observed statistics that, as it turns out, are heavily influenced by binary evolution.  Altogether, this alters our view about the most basic outcomes of massive-star mass loss --- are Wolf-Rayet stars and Type Ibc supernovae the products of single-star winds, or are they mostly the result of binary evolution and eruptive mass loss?  This is not fully settled, but mounting evidence points toward the latter.  This paradigm shift has far-reaching impact on other areas of astronomy, since it changes predictions for ionizing radiation and wind feedback from stellar populations, it may alter conclusions about star formation rates and initial mass functions in external galaxies, it affects the origin of various compact stellar remnants, and it determines how we use supernovae as probes of stellar evolution across cosmic time. ",
        "introduction": "Because of the scaling of luminosity with initial mass, relatively rare stars born with masses above $\\sim$10-20 $M_{\\odot}$ vastly outshine the much larger number of lower-mass stars.  The radiative output of massive stars is so intense that photon momentum can drive strong winds, and the energy transport through the stellar interior may jeaprodize the stability of the star itself.  Mass loss from massive stars has a {\\it deterministic} influence on the structure and evolution of those stars, which in turn have tremendous impact on other areas of astronomy.  Massive stars are the cosmic engines that provide most of the luminosity in star-forming galaxies.  Since massive stars have short lifetimes, their ultraviolet (UV) radiation (reprocessed in various forms) is the most observable tracer of star formation and is used to calculate the star formation rates \\citep{k98,ke12}.  Feedback in the form of UV radiation, stellar winds, and supernovae (SNe) stirs interstellar gas, driving turbulence and perhaps triggering new generations of stars by sweeping gas into dense filaments \\citep{el77}.  Through this feedback, massive stars have a profound impact on the evolution of disk galaxies \\citep{vdkf11}.  Massive star feedback may also terminate star formation locally, blowing giant bubbles \\citep{c05} and leaking hot processed gas into the Galactic Halo \\citep{putman12}.  Elemental yields from nuclear burning in the cores of massive stars and explosive nucleosynthesis in SNe provide the elements in the periodic table that pollute the interstellar medium (ISM), driving galactic chemical evolution and the metallicity ($Z$) evolution of the Universe.  Through their violent deaths as SNe, massive stars provide brilliant displays that permit us to dissect individual stellar interiors at distances of many Mpc while the star's inner layers are peeled away for us to see \\citep{filip97}.  They leave behind exotic corpses such as black holes, neutron stars, pulsars, magnetars, and all the bizarre high-energy phenomena that occur when compact objects remain bound in a binary system (e.g., \\citealt{rm06}).  Shock fronts in their beautifully complex SN remnants echo through the ISM for thousands of years after the bright SN display has faded \\citep{r08}.  Mass loss affects a star's luminosity, burning lifetime, apparent temperature, the hardness of its emitted radiation field, its He core mass, and it will profoundly impact the end fate of a star.  As a consequence, changes in estimates of mass-loss rates can alter expectations for their collective ionizing radiation, UV luminosity, winds, and SNe.  Thus, understanding massive stars and their mass loss also remains important as we push ever farther to the distant reaches of the Universe.  Inferences about the initial mass function (IMF; \\citealt{b10}) and variation of the star-formation rate through cosmic time \\citep{madau98,hb06} hinge upon converting reprocessed massive star UV luminosity to a collective star-formation rate. Gamma ray bursts (GRBs; \\citealt{g09,wb06}) and SNe also provide a probe of very distant populations and stellar evolution in early environments, but this requires us to understand the connection between stars of various types, their mass loss, and the eventual type of SN seen.  The first stars in $Z$-poor environments are expected to be very massive \\citep{bl04}, and hence, the scaling of mass loss with $Z$ affects abundances of the lowest $Z$ stars \\citep{bc05,by11} polluted by a small number of early SNe.  Since dense stellar winds can absorb portions of the Lyman continuum emitted by a hot star's photosphere \\citep{paco96}, estimates of mass-loss rates and their scaling with metallicity can profoundly impact topics as remote as reionization of the universe and the interpretation of spectra from galaxies at the highest redshifts currently being detected \\citep{lb01,f06,mw10}.  Although convenient recepies and simple scaling relations to account for the collective effects of mass loss and feedback will always be available, researchers working in these other branches of astronomy should not believe that such recepies are reliable to better than order-of-magnitude levels.  Even in the very local universe where we have excellent multiwavelength observations, there is still tremendous uncertainty in derived mass-loss rates for massive stars, and so there is large uncertainty in its influence on evolution.  This is exacerbated by the predicament that these uncertainties are largest for the most massive and most luminous stars, but these stars also tend to be the most influential.  Extrapolating to the early Universe is still quite risky, and should always raise eyebrows.  The aim of this review article is to provide a broad overview of the current understanding of mass loss and its influence on the evolution of massive stars, and to raise a flag of caution about the uncertainties involved.     \\subsection{The Importance of Mass Loss in Massive-Star Evolution}  For low- and intermediate-mass stars ($M_{ZAMS} < 8 M_{\\odot}$), wind mass loss is relatively unimportant for evolution until the final stages as the asymptotic giant branch (AGB) star transitions to a protoplanetary nebula.  For massive stars, however, mass loss cannot be ignored.  For most of their lives --- even on the main sequence --- massive stars above $\\sim$20 $M_{\\odot}$ shed mass in fast winds that affect their subsequent evolution (i.e. $\\int\\dot{M}dt$ can be a significant fraction of the stellar mass), and in post-main sequence phases the mass loss becomes critical in determining the type of resulting SN explosion.  Thus, mass loss is inexorably linked to evolution for massive stars. This article will not provide a detailed review of our understanding of stellar evolution from a theoretical perspective, since this has already been done in a number of excellent reviews. Although somewhat dated, \\citet{cm86} provide a good description of how models that incorporate mass loss compare to stellar evolution models without mass loss.  \\citet{mm00} provide a review of how the inclusion of rotation can influence single-star stellar evolution models, while \\citet{langer12} has reviewed more recent advances as well as important aspects of how close binarity may dramatically change the evolutionary paths of stars.  Finally, \\citet{whw02} discuss stellar evolution models with particular emphasis on the late pre-SN burning phases and their connection to core collapse.  Instead, this review will concentrate mostly on observational estimates of mass-loss and their impact on stellar evolution, because this is where the largest uncertainty currently resides.  Stellar evolution calculations must adopt prescriptions for mass-loss rates as input to their code, but these assumed rates determine the outcome of evolution.  Depending on precarious assumptions about wind strength, a red supergiant (RSG) can be made to evolve to the blue on the Hertzsprung-Russell (HR) Diagram, or not.  A massive star can be driven to the luminous blue variable (LBV) phase, or it can avoid it altogether. Moreover, the mass-loss rates typically used in stellar evolution models are time-averaged, even though eruptive and explosive mass loss events are observed.  Hence, there is great uncertainty in the predictions of all evolutionary models.  Tracks on the HR diagram are not plotted with error bars that reflect these uncertain assumptions.  Due to the need for time-averaged prescriptions in stellar evolution calculations, the steady line-driven winds of hot stars have consumed the majority of effort in understanding mass loss in the massive star community for the past three decades, both theoretically and observationally. There have been great leaps in quantitative non-LTE modeling of spectra influenced by winds.  A previous review by \\citet{kp00} (KP00 hereafter) dealt with line-driven winds of hot stars, focussing on the theory of the driving mechanism as well as common scaling relations used to convert observables to rates.  A more recent review by \\citet{puls08} also concentrated on the relatively steady line-driven winds of hot stars, providing a thorough discussion of wind theory and its connection to a wide array observational diagnostics.  Those reviews did not focus on more extreme winds of evolved massive stars or mass loss in binaries, and they did not discuss highly time-dependent mass loss associated with transient events (eruptions, explosions, mergers).  This review will therefore emphasize these latter topics, but it will discuss more recent developments, in particular the reduction in mass-loss rates and the consequent paradigm shift in stellar evolution.    \\subsection{Historical Perspective and Paradigm Shift}  Our present understanding of mass loss has undergone dramatic changes, due in large part to major shifts in our quantitative estimates of what mass-loss rates actually are.  The historical path to our current understanding has had some interesting swings.  Very early observations of transient events like Tycho's SN, P Cygni's 1600 AD eruption, and many nova eruptions, as well as the broad emission lines in Wolf-Rayet (WR) stars \\citep{wr1867} and the broad blueshifted absorption seen in spectra of many objects indicated the presence of outflowing material.  However, these were seen as rare stars or brief eruptive events, and the connection to the lives of normal stars was unclear.  Later, estimates for the solar wind \\citep{parker58} powered by hot gas pressure had such low mass-loss rates that winds would not matter much in stellar evolution.  There were some important early indications and expectations of mass loss from hot stars \\citep{sobolev60}, but the mass-loss rates were very uncertain.  Therefore, the dominant paradigm (e.g., \\citealt{pac66,pac67,pac71}) was that, as in low-mass stars, binary Roche-Lobe Overflow (RLOF) played a major role in making massive stripped-envelope stars like the He-rich WR stars.  The birth of UV astronomy triggered a revolution in our understanding of massive stars by providing the decisive evidence that all luminous hot stars have strong winds, indicated by deep P~Cyg absorption in their UV resonance lines \\citep{morton67}.  This led to a new paradigm wherein line-driven winds dominate the stripping of the H envelope, rather than binaries, leading to WR stars and Type Ibc SN progenitors via single-star evolution (i.e., the so-called ``Conti scenario''; \\citealt{conti76}).  As with UV observations, the development of IR detectors led to a similar recognition of the importance of dust-driven winds in red supergiants \\citep{gw71}.  With the assumption that mass loss in steady winds dominates mass loss, theorists could adopt simple prescriptions for that mass loss and calculate single-star evolutionary tracks on the HR Diagram (see reviews by \\citealt{cm86,mm00}).  These single-star models were able to provide a plausible explanation for observed distributions of stars, including the relative amounts of time spent in different evolutionary phases as O-type stars, WR stars, and RSGs. \\citet{massey03} has reviewed how these single-star models have been compared to observations and how observed statistics have been used to test and refine the single-star models.  This single-star paradigm then permitted one to extend even further, to compute models for the collective radiative output of entire stellar populations as a function of age (e.g., codes like {\\sc starburst99}; \\citealt{leitherer99}).     In the last decade this paradigm has shifted yet again because of two important realizations: 1) due to the effects of clumping (see below), empirical mass-loss rates for line-driven winds are lower than previously thought, and 2) other unsteady modes of mass loss are more important than previously recognized, due to increased estimates for mass ejected in episodic mass loss events \\citep{so06} and the very high binary fraction among massive stars \\citep{sana12}.  These changes to our view of mass loss are discussed in detail below. This may alter a huge number of predictions for the evolutionary tracks of stars and their variation with $Z$.  The pendulum appears to be swinging back to binaries as the dominant agent in massive star evolution, although not everyone is on the bandwagon.  Justifiably, this topic occupies much of the current debate and excitement in the field. ",
        "conclusions": "\\subsection{Take-home points}  We now have a fairly firm understanding of stellar winds of hot O-type stars \\citep{puls08}, relevant for most of their lives on the H-burning MS phase.  There remain substantial issues in understanding the physics of wind driving, magnetic fields, and angular momentum loss, and how these correspond to various observational diagnostics --- but in terms of the total mass-loss rates and their impact on stellar evolution, we know that wind mass loss is weaker than we used to think, due to the effect of clumping.  The consensus seems to be that mass-loss rates need to be reduced by a factor of at least 2--3 compared to rates derived observationally from standard $\\rho^2$ diagnostics assuming homogeneous winds \\citep{dj88,ndj90}.  Reduction by a factor of $\\sim$10 is probably an overstimate for most O-type stars, but not for later O-types and early B-type MS stars subject to the weak-wind problem.  Astronomers often scoff at a factor of 2 in any individual measurement, but we must recognize that systematically lowering $\\dot{M}$ by even a factor of 2 in models is very significant, and a factor of 3 is huge (current debate generally centers around a factor of 2 or 3).  A factor of only 2 lower $\\dot{M}$ would be like replacing mass-loss rates at $Z_{\\odot}$ with mass-loss rates currently used for 0.37 $Z_{\\odot}$ (i.e. lower than the LMC), whereas a factor of 3 would correspond to 0.2 $Z_{\\odot}$, similar to models for SMC stars (using $\\dot{M} \\propto Z^{0.69}$,\\citealt{vink01}). With moderatly weaker mass loss, we know that SMC models exhibit significantly different evolution than Milky Way stars.  Most importantly, they do not produce such a large population of H-free WR stars.  However, it is over this same metallicity range \\citep{massey03} that we are testing the validity of the physical ingredients in stellar evolution models.  This isn't good.  For post-MS supergiant phases and binary RLOF, the uncertainty in mass loss is far worse --- but these are also times when mass loss is strongest!  Given that single-star models neglect eruptive mass loss and binaries, but match many properties of the observed distributions anyway, predictions beyond the end of H burning are not unique. Extrapolations to low $Z$ and prescriptions for feedback from early stellar populations are therefore highly uncertain.  Researchers working in other fields where massive star feedback is relevant --- galaxy evolution, reionization, chemical evolution, low-metallicity stars, etc. --- should be cognizant of this.   \\subsection{Perspectives and Directions}  With such a high level of uncertainty in the most important phases of mass loss for massive stars, astronomers working on mass loss and massive star evolution have a few choices on how to proceed: 1) restrict their attention to MS stars, 2) work on Population III stars (see above), or 3) confront the complicated effects of binaries and time-dependent mass loss, as well as their influence on observed populations.  A major hurdle is disentangling the observed effects of single-star mass loss and binary RLOF.  On the one hand, Occam's razor encourages us to keep things simple, motivating us to understand single stars before we can hope to tackle the ``free parameter heaven'' of binaries.  On the other hand, the massive stars we observe are mostly binaries: A wise man once said that we should ``Make things as simple as possible, but not simpler''.  The message here is that ignoring binaries, but testing single-star models against observed statistical properties of massive stars (which are mostly binaries) leads us to misinterpret single stars too, because we are relying upon strong winds and rotation to compensate for actual outcomes of binary RLOF.   Since the observed statistics of massive-star populations are unavoidably contaminated by the effects of binary stars, one alternative approach for testing single-star models would be to place less emphasis on matching these observed {\\it statistics}, and more emphasis on matching individual observed stars with well-constrained physical parameters.  Masses for ``effectively single'' stars can be measured in wide binary systems where the stars will not interact (or O-type stars that have not yet interacted on the MS), and for these a detailed modern quantitative analysis can yield estimates of $\\dot{M}$, $T_{eff}$, $L$, rotational speed, abundances, etc.  The distance must be known, but the Magellanic Clouds and clusters with known distances can help, and {\\it GAIA} will soon lower distance uncertainties for the nearest massive stars in the Milky Way.  Doing this for a number of stars at various evolutionary stages and luminosities will build up a ``grid'' of observational constraints that models can aim to fit.  Much of this observational work has already been done.  Given the huge level of uncertainty in post-MS mass loss (for both RSGs and LBVs), another general approach for theorists could be to use stellar evolution codes as ``toy models'', to investigate the final outcome for a wide range of possible mass-loss prescriptions that include episodic mass loss.  For example, one could calculate evolutionary tracks for very massive stars that have lower wind $\\dot{M}$ during H burning, and then simply invoke one, two, or multiple sudden ejections of 10 $M_{\\odot}$ in post-MS phases to mimic LBV eruptions, in order to learn how the star's further evolution responds.  This approach may seem artificial, but the goal would be to reverse-engineer stars, not to provide uniquely correct evolutionary tracks.  In any case, much additional work on constraining post-MS mass loss of all forms is needed.  A number of specific possible future directions for theory and observations are listed below."
    },
    "1402/1402.3468_arXiv.txt": {
        "abstract": "Based on the Thomas-Fermi solution for compressed electron gas around a giant nucleus, we study electric pulsations of electron number-density, pressure and electric fields, which could be caused by an external perturbations acting on the nucleus or the electrons themselves. We numerically obtain the eigen-frequencies and eigen-functions for stationary pulsation modes that fulfill the boundary-value problem established by electron-number and energy-momentum conservation, equation of state, and Maxwell's equations, as well as physical boundary conditions, and assume the nucleons in $\\beta$-equilibrium at nuclear density. We particularly study the configuration of ultra-relativistic electrons with a large fraction contained within the nucleus. Such configurations can be realized for a giant nucleus or high external compression on the electrons. The lowest modes turn out to be heavily influenced by the relativistic plasma frequency induced by the positive charge background in the nucleus. Our results can be applied to heavy nuclei in the neutron star crust, as well as to the whole core of a neutron star. We discuss the possibility to apply our results to dynamic nuclei using the spectral method. ",
        "introduction": "The Thomas-Fermi model that was found independently by Thomas \\cite{thomas} and Fermi \\cite{fermi} in 1927 quantitatively describes neutral and ionized atoms of large electron-numbers with great success (see for example Refs.~\\cite{pomer1945}-\\cite{greiner82}). The Thomas-Fermi solution turns out to be exact when the electron-number goes to infinity \\cite{liebsimon}. Essentially, the Thomas-Fermi model is a semi-classical and mean-field approach to the problem of many electrons around a nucleus with a large number of protons. It describes a neutral or charged static equilibrium configuration of electrons around a nucleus with or without compression. While it turned out to be of limited use in the realm of atomic physics, it has been applied very successfully in astrophysical settings (see for example Refs.~\\cite{ruffinistella80,ruffinistella81,rotondo2,rotondo3,rotondo}).  In this article, on the basis of the Thomas-Fermi solution, namely the equilibrium configurations of electrons compressed around a giant nucleus, we investigate radial perturbations (electric pulsations) with spherical symmetry upon the equilibrium configurations. We find that the spectrum of pulsation modes is determined by two effects: (i) outside the nucleus the speed of sound of the electron gas determines propagation, with possible contributions from both non- and ultra-realtivistic zones, while (ii) inside the nucleus there is an additional contribution due to the relativistic plasma frequency induced by the nuclear positive charge background. For sufficiently low frequency modes this leads to the perturbation dying away exponentially within the volume of the nucleus, rendering it effectively unavailable for wave propagation. To study the configuration of ultra-relativistic electrons with a large fraction contained within the nucleus, we choose a proton number $Z=10^6$ for the purpose of practical numerical simulation and illustration of (ii). While the effects we observe are also present at smaller $Z\\approx 10^3-10^4$, a more realistic configuration that might be expected in the very deep crust of neutron stars in the form of \\textit{pasta equation of state}, they are less pronounced (see Figure \\ref{fig:10to2and4SpectPlot}), and high electron densities partially rely on the gravitational pressure in this case. Instead we choose $Z=10^6$ because here most electrons are kept inside the nucleus solely by the electric interaction, and $\\beta$-equilibrium is saturated throughout the nucleus. For the effects we observe it is essential that electron densities approach proton densities inside the nucleus, in any other case including high pressure laboratory setups, the spectrum of the vibrational modes would be dominated by the equation of state and the corresponding speed of sound of the electron gas, the only feature being a transition from non- to ultra-relativistic conditions (see discussion in the conclusions and Figure \\ref{fig:10to2and4SpectPlot}).  The electrons around a static nucleus are treated as a perfect fluid described by thermodynamic number-density $\\n$, energy-density $\\rh$ and pressure $\\p$ with non-vanishing electric potential and field. In addition to the equation of state at zero temperature, these physical quantities fully obey the Maxwell field-equation, Euler equation and the first thermodynamical law that follows from electron-number and energy-momentum conservations. This system is completely determined with appropriate physical boundary conditions. In order to study the perturbative electric pulsations, we have linearized these relations and equations, based on the prescription of Eulerian and Lagrangian perturbations of the Thomas-Fermi equilibrium configuration. As a result we obtain a homogeneous second-order differential equation for perturbations satisfying appropriate physical boundary conditions.  As a first step, we focus on the stationary solution $(\\propto e^{i\\omega t})$ with the characteristic eigen-frequencies $\\omega$ of electric perturbations (pulsations) of the Thomas-Fermi system, so as to understand what are time-scales (inverse frequencies) at which the system responds to external actions. In a future work we are planning to make use of these results to solve the problem of a general perturbation of the nucleus by means of the spectral method. Numerically and analytically solving this well-defined eigenvalue problem in a hybrid approach, we have obtained the eigen-frequencies and eigen-functions for electric perturbations of the Thomas-Fermi equilibrium configurations in spherical symmetry. Our study is an analogy of the classical investigation of stellar stability against gravitational pulsations \\cite{chandra,wheeler}, and also in the context of the Thomas-Fermi model for atoms similar calculations have been performed \\cite{greiner1,greiner2,wheeler3}. Our work is distinct because we consider a giant nucleus embedded in electron gas, and pay close attention to its interior, where the positive charge background of the nucleus induces a plasma frequency for electron oscillation. The purpose of this study is the application to astrophysical systems composed of nuclear matter and electrons, specifically to dynamical phenomena that could lead to creation of strong electric fields. We have drawn our attention to the ultra-relativistic limit corresponding to the polytropic equations of state ($\\p \\propto \\n^{\\Gamma_1}$) with adiabatic index $\\Gamma_1=4/3$. Our model could, however, easily be extended to include the transition to non-relativistic regimes, which are reached at compression radii one order of magnitude higher than the one ($r_\\mathrm{max}=1000\\,\\lambda_\\pi$) we are considering here (see Figure \\ref{fig:equilibriumPlot}). We present discussions on the physical relevance of our results and possible experimental testability in the lower crust and core of neutron stars.  The article is organized as follows: In Sec.~\\ref{basics}, we present fundamental equations, including the equation of state, describing the dynamics of an electron fluid in electromagnetic fields. In Sec.~\\ref{thomFer}, we show that in the static case these equations are reduced to the Thomas-Fermi equation, and present its solution. In Secs.~\\ref{pert} and \\ref{exSol}, we discuss the linear perturbations of the Thomas-Fermi solution, and derive the dynamical equations and boundary conditions that govern these perturbations, as well as obtain their stationary solutions by numerical and analytical methods. In the final Sec.~\\ref{conClu},  we present some discussions on the physical relevance of our results and the possible observational effects. We use the electron charge $-e$, natural units with $\\hbar = c = 1$ and Gauss convention for electrodynamics throughout the article, unless otherwise specified. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[scale=0.7]{multiPanel} \\caption{For $Z=10^6$, $\\Delta=2.85$, $r_\\mathrm{nuc}=285$, $r_\\mathrm{max}=1000$ and $s=28$ we show from top to bottom displacement $\\xi$, Eulerian (solid) and Lagrangian (dashed) number density perturbations $\\delta n$ and $\\Delta n$, equilibrium electric field (solid) and number density (dashed) $E_\\mathrm{eq}$ and $n_\\mathrm{eq}$, and Eulerian (solid) and Lagrangian (dashed) electric field perturbations $\\delta E$ and $\\Delta E$.} \\label{fig:multiPanel} \\end{figure}  Based on the approach of linear approximation, we quantitatively study the electric pulsation of the Thomas-Fermi equilibrium configurations around giant nuclear cores with fixed proton number-density $n_p$ and compression radius $r_\\mathrm{max}$. The eigen-frequencies $\\omega_s$ and eigen-functions $\\xi_s(r)$ of these electric perturbations are completely determined, except the absolute amplitude of $\\xi_s(r)$, corresponding to the arbitrary factor ${\\mathcal C}$ in Eq.~(\\ref{sol1}). The absolute amplitude of $\\xi_s(r)$ has in general to be determined by the space-time dependent strengths of external sources which initially trigger or continuously stimulate electric perturbations. In this case, the boundary-value problem turns out to be inhomogeneous in Eq.~(\\ref{pertEqn}) and/or boundary conditions (\\ref{pertOuterBound}). This is the subject of study in our upcoming work \\cite{hrx2014}, where we will use the spectrum obtained here to solve the inhomogenous problem by means of the spectral method. In the present article, we have completely determined the eigen-frequencies $\\omega_s$ by choosing $\\xi(r_\\mathrm{max})$-values in such a way that the linear approximation is valid for $\\omega_s\\xi_s(r)\\ll 1$. The amplitudes of eigen-functions $\\xi_s(r)$,  as well as the corresponding Lagrangian perturbations of thermodynamical quantities $\\Delta \\n$, $\\Delta \\rh$, $\\Delta \\p$ and electric field $\\Delta E$ are determined up to a constant factor, like ${\\mathcal C}$ in Eq.~(\\ref{sol1}).  Naturally the question arises whether the eigen-frequencies $\\omega_s$, in particular the lowest-lying one $\\omega_1$, can be experimentally tested, since these eigen-frequencies $\\omega_s$ are characteristic frequencies (time scales) of the Thomas-Fermi system responding to suitable external actions. To give an impression how the spectrum of our model behaves over a large range of compression radii $r_\\mathrm{max}$ we have integrated the full relativistic Eqs. (\\ref{fullrelFermi}) and (\\ref{eqn:Sturm-L}) for $Z=100$ and $Z=10^4$ with suiting approximate values for $\\beta$-equilibrium (see Table \\ref{tab:freq}) and show the results in Figure \\ref{fig:10to2and4SpectPlot}. It becomes clear from this figure that, while the transition from the non- to the ultra-relativistic equation of state is visible in the upper panel, which corresponds to common nuclei that could be probed in laboratory experiments, the constribution of the plasma frequency is negligible. In the lower panel, however, which corresponds to systems that might be expected in the lower crust of neutron stars, there is a visible effect of the plasma frequency on the spectrum at compression radii that are to be expected in such systems. In order to simplify the boundary-value problem and numerical calculation, we adopt a spherically symmetric model and find the eigen-frequencies and eigen-functions for spherical electric pulsations in the radial direction. Such spherical electric pulsations (the monopole vibrations) do not emit electromagnetic radiation (spin-one) due to the conservation of angular momentum. However, if structures in the lower crust such as the ones present in a \\textit{pasta equation of state}, or neutron star cores, do not possess the exact spherical symmetry, the dipole-component of electric pulsations appears, and this possibly leads to electromagnetic radiation with discrete frequencies. Such effects could be introduced to our spherically symmetric system by dropping the symmetry assumption for the perturbation, and expanding in the appropriate basis, e.g. spherical harmonics instead. In the context of the \\textit{pasta state} also other geometries such as axial and planar symmetries could be taken into consideration for the equilibrium state, with most of the calculation remaining unchanged. This procedure would allow to predict emission of electromagnetic radiation at frequencies to be determined by the details of the calculation (geometry, pressure), stimulated by dynamic events in the neutron star such as slowdown/spinup glitches or even more powerful events such as neutron star mergers. Furthermore, the effects we observe should influcence the dynamic compressibility of nuclear matter in the \\textit{pasta state}, but it is not clear when the tools will be available to probe such effects. If gravitational wave astronomy can be advanced to a precision that allows to probe tidal effects in neutron star mergers, it will allow to constrain crust equations of state.  \\begin{figure} \\centering \\includegraphics[scale=0.7]{10to2and4SpectPlot} \\caption{The eigen-frequencies $\\omega_{1-10}$ for $Z=100$ and $A=2 Z$ (upper panel), and for $Z=10^4$ and saturated $\\beta$-equilibrium (lower panel) versus the radius of compression. In the former case we can only observe the transition from non- to ultra-relativistic equation of state, in the latter case the influence of the nuclear plasma frequency becomes visible: for frequencies above $\\omega_\\mathrm{plas}$ the volume of the nucleus becomes available for wave propagation, rendering the slope less steep.} \\label{fig:10to2and4SpectPlot} \\end{figure}  The analytical pulsation frequencies obtained for the peculiar equilibrium configuration in Sec.~\\ref{exSol} could find their application in the context of neutral nuclear matter at/over nuclear density, for instance neutron star cores. The electrons and protons approximately have the same density profiles and occupy the same volume, except the very thin layer on the surface of neutron star cores \\cite{rotondo, rotondo2,rotondo3,rueda}, so the numerical results obtained for $Z=10^6$ can give insight into the pulsation behaviour near this surface. Moreover, if the amplitude of a perturbation is so large that the perturbation of the electric field is over the critical value $E_c=m_e^2c^3/(e^2\\hbar)$, electron-positron pairs are produced, leading to some observational consequences (see Ref.~\\cite{hrx2012, rx2013}).  \\begin{table} \\begin{tabular*}{\\textwidth}{lccc} \\hline $Z$ &\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ $10^2$\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ &\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ $10^4$\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ &$10^6$\\ \\\\ $A$ &$2\\times10^2$\\ &\\ \\ $2.2\\times 10^6$\\ \\ &\\ \\ $2.2\\times 10^8$\\ \\ \\\\ \\hline $r_\\mathrm{max}\\ \\ \\ \\ \\ (\\lambda_\\pi)$\\ \\ \\  & 8000 & 125 & 285\\\\ $\\langle n_e \\rangle\\ \\ \\ \\ \\ \\,(\\lambda_\\pi{}^{-3})$\\ \\ \\  & \\ \\ $5\\times 10^{-11}$ & $10^{-3}$ & $10^{-2}$\\\\ $\\omega_1$\\ \\ \\ \\ \\ \\ \\ \\,(MeV)\\ \\ \\ &\\ \\  $9\\times 10^{-3}$\\ \\  & 2.6\\ \\  & 5.3\\\\ $\\omega_\\mathrm{plas}$\\ \\ \\ \\ \\,(MeV)\\ \\ \\  & 13\\ \\  & 5.3\\ \\  & 5.3\\\\ $\\omega_\\mathrm{breath}$\\ \\ (MeV)\\ \\ \\  & 39\\ \\  & 11\\ \\  & 5.2\\\\ \\hline \\end{tabular*} \\caption{Comparison for suitable radius of compression $r_\\mathrm{max}$ of average electron density $\\langle n_e \\rangle$, corresponding fundamental mode $\\omega_1$, plasma frequency $\\omega_\\mathrm{plas}$ and nuclear breathing frequency $\\omega_\\mathrm{breath}$. Based on the assumptions of nuclear density, bulk nuclear incompressiblity of $K_0=220\\ \\mathrm{MeV}$, $Z=A/2$ for $Z=10^2$ and saturated $\\beta$-equilibrium in the other cases. The proton density $n_p$ in saturated $\\beta$-equilibrium equals $\\langle n_e \\rangle$ in the third column.} \\label{tab:freq} \\end{table}  To end this article, we briefly discuss the possibilities of external perturbations from nuclear cores, which probably cause the phenomenon of electronic pulsations discussed in this article. As an example, in Ref.~\\cite{wang}, the nuclear breathing modes of nuclear collective motion at the time scale of the nuclear force were discussed. We attempt to model the perturbation of nuclear cores and solve the inhomogeneous boundary-value problem to have further understanding of electric pulsations around nuclear matter \\cite{hrx2014}. For this pupose we make a rough estimate and consider only the bulk nuclear incompressibility $K_0\\approx 220\\,\\mathrm{MeV}$ as it dominates for large $A$. The nuclear breathing frequency is calculated according to \\cite{wang} as \\begin{align} \\hbar\\,\\omega_\\mathrm{breath}=\\sqrt{\\frac{\\hbar^2 K_0}{m_{_N} \\langle r^2 \\rangle}}=\\sqrt{\\frac{3 \\hbar^2 K_0}{5 m_{_N} r_\\mathrm{c}^2}} \\end{align} where $m_{_N}$ is the neutron mass, $r_c$ the radius of the nuclear core, and we have assumed homogenous distribution of nuclear matter in the core. In Table \\ref{tab:freq} we compare the nuclear breathing modes to the plasma frequencies. While the former depend on the size of the nucleus, the latter depend only on proton density, and we observe that they are of comparable size for $Z\\approx 10^6$. As we have outlined in \\cite{zeldoProc} a driving force proportional to the radius, corresponding to a homogeneous nuclear breathing mode, can easily be modelled using the spectral method and the inner product (\\ref{innerProd}), effectively calculating time-dependent coefficients for the pulsation modes derived in this article.  It is worthwhile to mention that in gravitational collapses of macroscopic cores at/over nuclear density, strong dynamical variations of collapsing cores can be induced at or over the rate of nuclear interactions, which cannot be treated as linear perturbations, and these variation are no longer stationary. Non-perturbative calculations showed that this results in strong variations of the electric field, leading to the formation of electron-positron pairs and photons at high energy- and number-densities \\cite{hrx2012,rx2013}. Further quantitative studies on this issue are being conducted."
    },
    "1402/1402.6551_arXiv.txt": {
        "abstract": "We report a time-series analysis of satellite photometry of the non-radially pulsating Oe star $\\zeta$~Oph, principally using data from \\emph{SMEI} obtained 2003--2008, but augmented with \\emph{MOST} and \\emph{WIRE} results. Amplitudes of the strongest photometric signals, at 5.18, 2.96, and 2.67~\\pd, each vary independently over the $5\\nicefrac{1}{2}$-year monitoring period (from $\\sim$30 to $\\lesssim$2~mmag at 5.18~\\pd), on timescales of hundreds of days. Signals at 7.19~\\pd\\ and 5.18~\\pd\\ have persisted (or recurred) for around two decades.  Supplementary spectro\\-scopic observations show an H$\\alpha$ emission episode in 2006; this coincided with small increases in amplitudes of the three strongest photometric signals. ",
        "introduction": "\\label{sec:intro}     $\\zeta$~Oph\\footnote{HD 149757} is the nearest O-type star, and one of the brightest ($\\pi = 8.91\\pm0.20$~mas, \\citealt{vanLeeuwen07}; O9.5$\\,$Vnn, $V \\simeq 2.6$, \\citealt{Sota11}). Detailed spectro\\-scopic investigations of velocity-resolved absorption-line structure have been facilitated by its brightness and exceptionally rapid rotation ($v_{\\rm e}\\sin{i} \\gtrsim 400$~\\kms; \\citealt{Howarth01}, \\citealt{Villamariz05}); periodic line-profile variability, discovered by \\citet{Walker79}, has subsequently been widely interpreted in terms of non-radial pulsations (NRP; \\citealt{Vogt83}, \\citealt{Reid93}, \\citealt{Kambe97}).  H$\\alpha$ emission episodes lasting, typically, several weeks have been observed to occur every few years (e.g., \\citealt{Ebbets81, Kambe93}); the inferred circumstellar decretion disk is probably causally associated with rapid rotation, and possibly with NRP \\citep{Cranmer09}, though the latter remains a open \\mbox{issue}.  Spectro\\-scopic line-profile variability associated with NRP is primarily sensitive to sectoral pulsation modes (since tesseral modes readily lead to cancellation in velocity space). Moreover, the requirements of high signal-to-noise ratio and high resolution typically limit spectro\\-scopic time series to only a few nights, resulting in further detection biases, towards short periods and large amplitudes. These observational constraints contrast with satellite-based photo\\-metry, which can yield precise measurements over an extended time period, thereby affording the opportunity to investigate pulsation characteristics in a parameter space inaccessible to spectro\\-scopic study.  For $\\zeta$~Oph, this opportunity was exploited by \\citet{Walker05}, who found a number of periodic signals in 24~days of high-cadence, near-continuous photo\\-metry from the \\emph{Microvariability and Oscillations of STars (MOST)} satellite.  Here we report new results from the \\emph{Solar Mass Ejection Imager (SMEI)} and \\emph{Wide-field Infra\\-Red Explorer (WIRE)} missions; for completeness, we also include our independent re-analysis of the \\emph{MOST} data. ",
        "conclusions": "\\label{sec:conclusion}  Data from the \\emph{SMEI, WIRE, \\emph{and} MOST} satellites have been analysed to investigate periodic signals in broad-band optical photometry of the Oe star $\\zeta$~Oph obtained over a span of almost 6~years.  We confirm multiperiodic variability at the $\\sim$10~mmag level and, for the first time, track systematic changes in signal amplitudes on timescales of order $\\sim{10}^2$~d; some signals, while not continuously detectable, are never\\-the\\-less present in observations separated by 20~years. There is tentative evidence of a photometric signature of the 2006 emission-line episode; although no direct correspondence is evident between overall photometric and spectro\\-scopic activity, this \\emph{may} reflect different roles of radial and non-radial modes in the formation of a decretion disk."
    },
    "1402/1402.4762_arXiv.txt": {
        "abstract": "{The All Sky Camera (ASC) was developed as an universal device for a monitoring of the night sky quality and night sky background measurement. % ASC system consists of an astronomical CCD camera, a fish eye lens, a control computer and associated electronics. The measurement is carried out during astronomical twilight. The analysis results are the cloud fraction (the percentage of the sky covered by clouds), night sky brightness (in $mag/arcsec^{2}$) and light background in the field of view of the camera. The analysis of the cloud fraction is based on the astrometry (comparison to catalogue positions) of the observed stars.} ",
        "introduction": "An important issue of the night sky monitoring for High-Energy Astroparticle Detectors (HEAD) is the interaction of the monitoring instrument with the measurement of high energy cosmic ray particles. Contrary to lidars or radars that are active devices interfering at least to some extend with the HEAD data taking, the all sky cameras together with robotic telescopes or pyrometers are passive instruments that can work completely independently on HEAD. The benefit of all sky cameras is that the device works during the whole measurement period, it monitors atmosphere inside full HEAD Field Of View (FOV) and it does not influence the observation of high energy atmospheric showers.  All sky cameras that are simple and cheap devices can be used to characterize immediate cloudiness, night sky brightness and their time variability not only at running experiments but also at candidate sites of future observatories \\cite{bib:ICRCcam}. On the other hand compared e.g. to lidars that deliver full atmospheric profiles of measured quantities (cloud heights, aerosol optical depth vs height), the all sky cameras cannot provide such a complex and detailed information.  This contribution focuses on basic hardware and data analysis principles related to all sky cameras. ASCs developed by Joint Laboratory of Optics of Palacky University and Institute of Physics ASCR in Olomouc are taken as an example.    \\subsection{System Description}  The basis of the system is an astronomical camera G$1-2000$ (Moravian Instruments a.s., Czech Repu\\-blic, www.mii.cz), which is equipped with a CCD chip ICX$274$AL by SONY company. The quantum efficiency of the CCD chip is higher than $50 \\%$  at $450$--$550$ nm and its spectral sensitivity covers the range from approximately $400$ to $900$ nm. The electronics has low read-out noise, a 16-bit analog-digital converter and a resolution of $1600\\times1200$ pixels. The camera is equipped with a fish-eye varifocal lens (the field of view is $185$ degrees) and an electronic controlled iris. This setup is capable of detecting a star with visual magnitude up to 6 mag in zenith. %  A miniPC computer with USB I/O controls the system and processes the data. The I/O card controls the iris and the power switch of the camera. All electronics and a camera body (see Figure \\ref{electronics}) are weatherproof. The system power could be supplied by a solar power system or local power line. The switching electronics turns the system ON after sunset and keeps it going during astronomical night. The system is OFF during the day (to save energy). Temperature of the electronics is controlled and stabilized using internal heating system during winter time to protect the ASC system.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.9\\textwidth]{atmohead2013-Mandat-electronics.pdf} \\caption{Electronics and the inside of the camera box.} \\label{electronics} \\end{center} \\end{figure*}  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.43\\textwidth]{atmohead2013-Mandat-night_sky1.pdf} \\hspace{0.5cm} \\includegraphics[width=0.43\\textwidth]{atmohead2013-Mandat-night_sky_clouds1.pdf} \\caption{Left: Image of clear night sky. Right: Partly cloudy night skies. Moon and clouds are visible close to horizon.} \\label{nightSky} \\end{center} \\end{figure*}     \\subsection{Calibration of the field of view}  The fish-eye varifocal lens has $185^\\circ$  FOV and its aberration distorts the image. The correct transformation of the incident angle to the imaged pixel position is very important. The mechanism of the calibration is as follows: the light source (a spot as small as possible) is placed at a defined distance (at least 10 m), then the body of camera is rotated around the axis such that the incident light angle (corresponding to the FOV angle) varies between 0 and 90 degrees. The position of the light spot is analyzed for each angular step. This procedure is repeated for different plane cuts of the lens FOV. The final data are analyzed and fitted with a polynomial function. This calibration is regularly checked on-site by comparing the position of detected stars with the stars in a catalog. On-site measurements are also used to determine the variation of the sensitivity of the system with the zenith angle. \\\\  \\subsection{Night sky data taking and principles of analysis} The ASC system takes full sky images  periodically in predefined moments (time period is limited by $2\\times$ exposure time - night sky image and dark image). Examples of clear and partly cloudy night sky are shown in Figure \\ref{nightSky}. %     The measurement process consists of the following steps. Initially, every exposure of a night sky is followed by a dark frame image so that subtraction of one from the other enables significant reduction of noise in the image. The horizontal coordinates of each image pixel are then calculated with the calibration data and calibration constants used to determine the horizontal coordinates, azimuth and elevation of each image pixel. The total charge of each pixel is proportional to the light intensity, but the intensity of the final image varies with respect to zenith angle due to vignetting and distortion of the lens. The calibration of this zenith angle intensity dependence is obtained by considering extinction in air masses, optical distortions and vignetting variation as a function of zenith angle. The resulting attenuation at, for example, $60$ degrees zenith angle corresponds to $4$ mag (see Fig.~\\ref{zen_mag}).  To estimate the cloud fraction, we investigate the presence of stars in the field of view of the camera. We separate the full sky to approximately 70 segments and analyze each segment of the sky individually to obtain the cloudiness. As a star catalog we use Yale Bright Star Catalog BSC$5$ \\cite{bib:starcat}, particularly the visual magnitudes of the stars. First we estimate the position of each segment on the sky using the data from the calibration. This information is then further improved using the brightest stars from the catalog in each segment. Once the position of a segment on the sky is fixed, the stars for the segment are read from the catalog and selected according to the expected sensitivity at the given zenith angle. For every star in the catalog, we look for a detected star within the angular limit of $1^\\circ$. If a star that fits this criteria is found, then the catalog and detected stars are flagged as paired. Figure \\ref{cloudiness_analysis} shows a result from cloudiness analysis of partly cloudy night sky. The ratio of paired / unpaired stars (for all segments) gives final cloudiness of the night sky. Typically, we check about 500 catalog stars for the whole observed sky (up to $60^{\\circ}$ due to high distortion close to horizon). The error of the algorithm was calculated using artificial cloud simulations (Fig.~\\ref{error}) and the total uncertainties of the cloudiness calculation are 2\\% for cloudiness $<$20\\% and 5\\% for cloudiness $>$80\\%. The cloudiness of the night sky is computed up to the maximal zenith angle with 10$^{\\circ}$ steps.     Brightness of the night sky could be evaluated directly from a light flux. The system has been calibrated using a darkness sky measurement tool Unihedron SQM-LE \\cite{bib:sqm}. The night sky was scanned overnight and recorded data both from the camera and SQM instrument were compared to each other. Subsequently the conversion formula is obtained. The data represent the brightness of the night sky in $mag/arcsec^2$ and in visible light spectrum.    \\vspace{0.5cm} \\subsection{Limits of the cloudiness analysis}  The measurement using the ASC is limited by local weather conditions, camera misting and Moon position. The local weather limitations mean for example rain or snow, fog, haze and water condensation on the camera. Examples of rain, fog and snow are shown in Figure \\ref{fog}. The algorithm usually fails if the lens of the ASC is obscured in some way, for example by snow, water resulting from rain fall, haze or fog surrounding ASC, or even bird droppings from birds found locally. The presence of the Moon saturates the image and increases the background in the image, and consequently the algorithm for cloud detection could not be used at such times. ",
        "conclusions": "The ASC, its application and limits are presented. The result of the analysis is the night sky cloud cover and night sky brightness. The main advantage of the all-sky camera is that it is a fully autonomous and self-reliant tool which can be used to monitor the night sky in even the world's most remote places, which happen to be just the ones most promising for astronomically related applications. The autonomy of the camera significantly contributes to the extremely low cost for its operation. The costs involved are concentrated mainly in the deployment phase, namely the costs of the parts and assembly of the camera, shipping (varies depending on destination, the shippable weight being 10 kg) and installation (two people working for two or three days). As demonstrated in \\cite{bib:ICRCcam} the camera can operate without maintenance for months and years after installation. However, it is advantageous to have local support for the case of unexpected external hindrances, but in most cases, a nonspecialist local personnel is sufficient to fix most of the possible issues.     \\begin{figure}[t] \\centering \\includegraphics[width=0.44\\textwidth,height=0.175\\textheight]{atmohead2013-Mandat-zenith_mag.pdf} \\caption{Zenith angle to magnitude dependence. The red points show sensitivity thresholds % in different azimuthal angles as a function of the zenith angle. The blue line is the geometric calibration - distortion description. } \\label{zen_mag} \\end{figure} \\begin{figure}[t] \\centering \\includegraphics[width=0.45\\textwidth,height=0.26\\textwidth]{atmohead2013-Mandat-error.pdf} \\caption{Error of the algorithm calculated using artificial cloud simulations with cloud fractions from $0\\%$ - $100\\%$ with step of $5\\%$.  Horizontal axis shows the simulated cloud cover and vertical axis corresponds to values calculated by the algorithm (blue dots). The error bars (red) represent the statistical estimation of the algorithm uncertainty.} \\label{error} \\end{figure} \\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{atmohead2013-Mandat-analysis1.pdf} \\caption{An example of analysis of partly cloudy night sky. The image on the left shows RAW image of the sky. The image on the right shows the analysis results, the yellow crosses indicate the detected stars, green circles catalog stars and red circles represent catalog stars without detected pairs - the region covered by cloudiness. The red circle indicates the limiting zenith angle ($60^{\\circ}$).} \\label{cloudiness_analysis} \\end{figure*}   \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=0.45\\textwidth]{atmohead2013-Mandat-fog1.pdf} \\hspace{0.5cm} \\includegraphics[width=0.45\\textwidth]{atmohead2013-Mandat-water1.pdf} \\caption{Left: Fog or haze affects the visibility of night sky, the nearby tower is not visible (the visibility fades with the fog/haze). Right: Lens of the ASC covered with water drops. } \\label{fog} \\end{center} \\end{figure*}     \\vspace{2cm}  \\footnotesize{{\\bf Acknowledgment: }{This work was supported by the  Ministry of Education Youth and Sports of the Czech Republic under the projects LE13012, 7AMB12AR013 and LG13007}}"
    },
    "1402/1402.5975_arXiv.txt": {
        "abstract": "A population of intermediate-mass black holes (BHs) is predicted to be freely floating in the Milky Way (MW) halo, due to gravitational wave recoil, ejection from triple BH systems, or tidal stripping in the dwarf galaxies that merged to make the MW. As these BHs traverse the gaseous MW disk, a bow shock forms, producing detectable radio and mm/sub-mm synchrotron emission from accelerated electrons. We calculate the synchrotron flux to be $\\sim \\rm 0.01-10\\, mJy$ at GHz frequency, detectable by {\\it Jansky Very Large Array}, and $\\sim 10-100\\,\\mu\\rm Jy$ at $\\sim10^{10}-10^{12} \\,\\rm Hz$ frequencies, detectable by {\\it Atacama Large Millimeter/sub-millimter Array}. The discovery of the floating BH population will provide insights on the formation and merger history of the MW as well as on the evolution of massive BHs in the early Universe. ",
        "introduction": "Galaxies grow through accretion and hierarchical mergers. During the final phase of the merger of two central black holes, anisotropic emission of gravitational waves (GW) kicks the BH remnant with a velocity up to a few hundreds $\\kms$ \\citep{baker2006, campanelli2007, blecha2008}. Additionally, BHs can be ejected from triple systems \\citep{kulkarni2012, hoffman2007}, or result from tidally-stripped cores of dwarf galaxies \\citep{bellovary2010}. For GW recoils, the typical kick velocity is large enough for the BHs to escape the shallow gravitational potential of low-mass galaxies, but smaller than the escape velocity of the MW halo. This is also the case for triple systems as long as the kick velocity is $<500 \\,\\kms$. Consequently, a population of floating BHs formed from mergers of low-mass galaxies are trapped in the region that eventually makes the present-day MW \\citep{madau2004, volonteri2005, libeskind2006}. Previous studies suggested that more than $\\sim$ 100 floating BHs should be in the halo today, based on a large statistical sample of possible merger tree histories for the MW halo today \\citep{oleary2009, oleary2012}. This population of recoiled BHs is supplemented by BHs from tidally disrupted satellites of the MW \\citep{bellovary2010}. The discovery of this BH population will provide constraints on the formation and merger history of the MW as well as the dynamical evolution of massive BHs in the early Universe.  It has been proposed that a compact cluster of old stars from the original host galaxies is carried by each floating BH \\citep{oleary2009}. In this Letter, we propose an additional observational signature of floating BHs, using the MW gas disk as a detector. As the BHs pass through the MW disk supersonically they generate a bow shock, which results in synchrotron radiation detectable at radio and mm/sub-mm frequencies.  The paper is organized as follows. In \\S~\\ref{sec:interact}, we discuss the interaction between BHs and the gas in the MW disk. In \\S~\\ref{sec:obs}, we calculate the synchrotron radiation from the bow shocks produced as the BHs cross the MW disk, and discuss the detectability of this radiation. Finally, in \\S~\\ref{sec:sum}, we summarize our results and discuss their implications. ",
        "conclusions": "\\label{sec:sum} If a floating BH happens to pass through the MW disk, then the non-thermal emission from the accelerated electrons in the bow shock around the BH should produce detectable signals in the radio and mm/sub-mm bands. The radio flux $\\sim 0.01-10\\;\\rm mJy$ is detectable by JVLA, while the mm/sub-mm flux $\\sim 10-100\\,\\mu\\rm Jy$ is detectable by ALMA.  The density distribution of floating BHs in the MW has been studied by O'Leary \\& Loeb (2009, 2012) and by Rashkov \\& Madau (2014). High resolution simulations show that there is a BH of mass $\\sim2\\times10^5\\msun$ within a few kpc from the Galactic center \\citep{rashkov2014}.  Observations of the Galactic disk can be used to infer $\\n0$ and $\\t4$. The BH speed $\\vbh$ can then be estimated from the Mach cone angle. The maximum Lorentz factor $\\gmax$ can be inferred from the peak of the synchrotron spectrum. This, in turn, yields $\\b-5$ based on Eq.(3). From the slope of the synchrotron spectrum, the power law index $p$ can be estimated. Finally, with the above parameters constrained, the synchrotron flux can be used to calibrate $\\mbh$. The above interpretation can be verified by observing the properties of the star cluster carried by the floating BHs \\citep{oleary2009, oleary2012}. The diffuse X-ray emission from the BH and synchrotron emission from the bow shock is supplemented by stellar emission from the star cluster around it. Since the total mass of the star cluster is much smaller than $\\mbh$, gravity is dominated by the BH, and thus the stars do not effect the bow shock. One can measure $\\mbh$ spectroscopically from the velocity dispersion of the stars as a function of distance from the BH, and verify consistency with the synchrotron flux estimate."
    },
    "1402/1402.5528_arXiv.txt": {
        "abstract": "The galaxy pair NGC~5580 and NGC~5588 are part of a loose group of galaxies. They are surrounded by steep-spectrum, extended radio emission which was previously suggested to be a down-scaled example of Mpc-size radio haloes present in galaxies clusters.  We present a multi-frequency study of the radio-emission aimed to clarify its nature. The source has been observed with the Giant Meterwave Radio Telescope (\\gmrt{}) at 235, 325 and 610 MHz and the images obtained were combined with archival data to cover the frequency range 150--1400 MHz.  The new observations revealed the presence of a second, fainter lobe on the South-East of NGC~5580. The spectral index study of the source shows a flattening of the spectrum (which implies a younger particle population) close to the two galaxies. We argue that the extended radio emission is the remnant of a past activity cycle of the active galactic nucleus (AGN) present in NGC~5580 and therefore a notable example of a dying radio galaxy located outside a dense environments. ",
        "introduction": "The galaxy pair NGC~5580/NGC~5588 is surrounded by an unidentified, bright and steep-spectrum radio emission. These two galaxies form a loose group with NGC~5557 \\citep[LGG~378;][]{Garcia1993}. Other publications define the group to include up to 13 galaxies \\citep[e.g., G~141;][]{Geller1983}, while \\cite{Fouque1992} group NGC~5590 with a few more galaxies (NGC~5557, NGC~5529 and UGC~9113; R7-A6) but exclude NGC~5588. The peculiar radio emission covering (and apparently connecting) the two galaxies has been so far ignored. Only \\cite{Delain2006} noticed it in the WENSS survey and classified it as a ``galaxy group radio-halo''. However, the source is not an X-ray emitter ($\\log L_X < 41.2$ erg/s) and the galaxies are part of a rather small ($\\lesssim 13$ galaxies) and loose group with high velocity dispersion \\citep[461 km/s;][]{VanDriel2001}. Here we present pointed radio observations aimed at understanding the nature of this emission which may be a down-sized version of the Mpc-scale radio halos present in massive clusters of galaxies \\citep{Venturi2008, Bonafede2009, Feretti2012} or it may be related to past episodes of AGN activity \\citep{Schoenmakers2000, Parma2007, Murgia2011}.  The active stage of an AGN can last several tens of Myrs. During this period the radio emission associated with the galaxy is powered by a continuous flow of energy from the AGN which reached distances several times the hosting galaxy size via plasma beams and jets. In galaxy clusters the injection of energy sustains the growth of the radio lobes against the pressure in the hot X-ray emitting external medium into which they expand. After this period the AGN activity ceases or is dramatically reduced to the point that the plasma outflow can no longer sustain the radio source, which undergoes a period of fading (dying phase). During this phase only the radio lobes remain visible. After a time that depends on the source environment, also the emission becomes too faint to be detected and the source vanishes.  Dying radio galaxies are an important but poorly explored stage of AGN evolution. It is believed that the low-frequency radio emission from the fading radio lobes lasts longer \\citep[$\\sim 10^8$ yr;][]{Murgia2011} if their expansion is prevented by the pressure of a dense intra-cluster medium (ICM). In this way cosmic ray electrons (CRe) lose their energy only through synchrotron radiation and inverse Compton scattering, while expansion plays only a marginal role. It is also possible for the source to restart before the emission from a previous burst faded completely \\citep{Schoenmakers2000}. The best example of fossil radio lobes seen in a currently active galaxy is 3C 338 \\citep[see e.g.,][]{Gentile2007} where steep-spectrum lobes of the source are clearly disconnected from the currently active jets.   In this paper we study the radio emission surrounding the galaxy pair NGC~5580/NGC~5588. In Sec.~\\ref{sec:obs} we discuss the set-up and the data reduction of the \\gmrt{} observations, and the archival \\vla{} data used. In Sec.~\\ref{sec:images} and Sec.~\\ref{sec:spidx} we analyse the images and the spectral behaviour of the source. In Sec.~\\ref{sec:discussion} we discuss various explanations for the radio source.  In this paper we adopt a $\\Lambda$CDM cosmology with $H_0 = 71$, $\\Omega_M = 0.27$, $\\Omega_{\\Lambda} = 0.73$. NGC~5580 is at a redshift $z = 0.0107$ ($D_L=45.6$ Mpc), which gives a scale of 13 kpc/arcmin. ",
        "conclusions": "\\label{sec:concludions}  Using the \\gmrt{} we observed the extended radio emission located in the galaxy group hosting NGC~5580 and NGC~5588. The radio emission is peculiar from, both, a morphological and spectral point of view. It shows one large \\textit{bubble} with no apparent internal structure, encompassing the lenticular galaxy NGC~5580 and the spiral galaxy NGC~5588. The two galaxies sit at the edges of this extended emission, which has been previously classified as an uncommon example of radio halo in a galaxy group \\citep{Delain2006}. The new observations revealed unambiguously the presence of a second \\textit{bubble} of similar size as the first one, located to the South-East of NGC~5580. The spectral shape of the source shows a flattening of the SED around both galaxies.  The nature of this emission is probably related to a single past episode of AGN activity from NGC~5580, which would be able to explain both the source morphology and spectrum. In this case, the presence of NGC~5588 would be a coincidence and NGC~5580 would be a notable case of dead radio galaxy located outside a dense cluster. This implies that AGN do not require dense environments to undergo recurrent episodes of activity. If this picture will be confirmed by other examples, this would support a scheme where AGN are not preferentially triggered by mergers, in line with the recent findings of \\cite{Karouzos2013}.  The fact that these type of sources are not often found outside of galaxy clusters may not be related to the intrinsic source population, since many radio sources are isolated and each of them must enter a dying phase. Indeed, \\cite{Murgia2011} suggested that the probability to observe an isolated dying source is comparatively low because radio sources in under-dense environments may have a slower duty-cycle than those in dense environments and/or their lobes may fade faster given the poor confinement of the radio-emitting plasma. However, the observation of the fading phase of NGC~5580 shows that dying radio galaxies can be studied also in relatively isolated objects. Future large-scale low-frequency radio surveys, as those being carried on by the Low Frequency Array (LOFAR) and in the future by the Square Kilometre Array (SKA), will provide enough statistics to shed light on this intriguing topic."
    },
    "1402/1402.1551_arXiv.txt": {
        "abstract": "Over the past decade it has become apparent that a class of `bursting pulsars'  exist with the discovery of PSR~J1752+2359 and PSR~J1938+2213. In these pulsars, a sharp increase in the emission is observed that then tends to systematically drop-off from pulse-to-pulse. In this paper we describe the discovery of such a relationship in high-sensitivity observations of the young (characteristic age of $90,000$~yrs) 0.33~s pulsar B0611+22 at both 327~MHz and 1400~MHz with the Arecibo radio telescope. While Nowakowski previously showed that B0611+22 has mode-switching properties, the data presented here show, for the first time, that this pulsar emits bursts with characteristic time-scales of several hundred seconds. At 327~MHz, the pulsar shows steady behaviour in one emission mode which is enhanced by bursting emission slightly offset in pulse phase from this steady emission. Contrastingly at 1400~MHz, the two modes appear to behave in a competing operation while still offset in phase. Using a fluctuation spectrum analysis, we also investigate each mode independently for sub-pulse drifting.  Neither emission mode (i.e.~during bursts or persistent emission) shows the presence of the drifting sub-pulse phenomenon. The bursting phenomena seen here appears to be a hybrid between bursting seen in other pulsars and the bistable profile illumination behaviour reported in two other pulsars by Rankin et al. Further examples of this cross-frequency behaviour are required, as this phenomenon may be quite common among the pulsar population. ",
        "introduction": "High-precision timing observations which probe neutron star physics, general relativity, and other phenomena rely on tracking rotational phase as measured by template matching of stable integrated pulse profiles \\citep[for a discussion, see][]{Lorimer:2005fk}. Upon closer examination, however, the radio emission displays a rich variety of examples in which pulsars appear to flip between various states on a wide range of timescales.  Shortly after the discovery of pulsars, in a remarkable series of papers making use of the Arecibo telescope, \\cite{1970Natur.227..692B, 1970Natur.228...42B,1970Natur.228..752B,1970Natur.228.1297B} showed that pulsars exhibit nulling, mode-changing and sub-pulse drifting phenomena in their pulse-to-pulse emission. Nulling, the abrupt halting of the emission mechanism for a few or many hundreds of pulses \\citep{1970Natur.228...42B} can, be viewed as an extreme version of mode-changing \\citep{1970Natur.227..692B,1970Natur.228..752B,1970Natur.228.1297B} during which the pulse profile exhibits two or more stable pulse shapes. Although not universal, an analysis of pulse nulling across the population of isolated pulsars by \\cite{2007MNRAS.377.1383W} tends to suggest that the fraction of time spent in the null state increases with characteristic age. Sub-pulse drifting \\citep{1968Natur.220..231D,1973NPhS..243...77B}, where clearly ordered subpulses drift across the main pulse window at a fixed rate characterized by a subpulse spacing and a repetition period appeared to be preserved over an interval of time spent in the null state \\citep{1978MNRAS.182..711U}.   With the passage of time, more discoveries have blended and questioned these classifications. \\cite{1970ApJ...162..727H} discovered the first of a number of pulsars that appear to combine mode changing with sub-pulse drift, where these pulsars generally assume one of several sub-pulse drifting modes.  Also, \\cite{MNR:MNR10512} have shown that some pulsars exhibit what they term ``bistable profile illumination'', which is single-pulse drifting that mimics the abrupt beginnings of mode-changing but then gradually drifts into phase with the normal mode. \\cite{2007MNRAS.377.1383W} suggest that nulling and mode changing are different manifestations of the same phenomenon.  Other extraordinary pulsars that may relate nulling to mode-changing occur on a wide range of timescales. A few of these discoveries have been the rotating radio transients (RRATs; \\cite{2006Natur.439..817M}), the intermittent pulsars \\citep{Kramer28042006, 0004-637X-746-1-63, 2012ApJ...758..141L} and the observation of bursting emission, which was until now only seen in J1752$+$2359  \\citep{0004-637X-600-2-905} and J1938$+$2213 \\citep{2003PhDT.........2C, Lorimer2013}. The bursting pulsars bear some resemblance to conventional nulling pulsars, with the exception that the pulse to pulse intensities are modulated by a decaying envelope before abruptly returning to a null or faint state. While the RRATs emit only occasional individual pulses spaced anywhere between seconds and years, intermittent pulsars switch on and remain steady before switching off again over timescales of days to years. Over these longer timescales, where it is possible to measure the spin period derivative precisely, it is observed that the ``on'' and ``off'' states for the intermittent pulsars are associated with two different spin-down rates\\citep{Kramer28042006}.  A similar behaviour was recently seen in 17 non-intermittent pulsars \\citep{DataHome} which switch between different spin-down states in a quasi-periodic or even chaotic fashion \\citep{2013MNRAS.428..983S}, and are sometimes accompanied by correlated pulse profile changes. \\cite{DataHome} suggest that the mixture of spin-down states could even account for the ``timing noise'' variations seen in many pulsars and suggest that most, if not all, normal pulsars have multiple spin-down states indicating global changes in magnetospheric current densities. One implication of these results is that mode changing (and perhaps, by association, bursting) could be connected to spin-down rate changes on much shorter timescales. Very recently, synchronous X-ray and radio state switching was reported in the classical mode-changing pulsar B0943+10~\\citep{2013Sci...339..436H}. In this case, non-thermal unpulsed X-rays are observed during the bright radio emitting phases, while an additional pulse thermal component is present during the radio-quiet phase. A better understanding of the observational phenomenology involved is now required in order to make further progress.   In this paper, we present a previously unseen bursting emission phenomenon in PSR~B0611$+$22 which was discovered by \\cite{DAVIES:1972fk}. This young pulsar with a characteristic age of $90,000$~years, has considerably high timing noise \\citep{1980ApJ...237..206H}. This high level of noise could indicate that there is still some underlying phenomena that is not being accounted for in the timing models. At first glance, it appears to be a normal pulsar with a simple pulse shape - essentially Gaussian in appearance that is approximately 80\\% linearly polarized \\citep{1989ApJ...346..869R}. Yet, it has been shown that the integrated pulse shape varies at different times \\citep{1980ApJ...239..310F}. \\cite{1992msem.coll..280N} has also shown that there is a correlation with the brighter pulses and their phase alignment, and later claimed that the timing noise was due to mode switching \\citep{2000AAS...19713004N}. Other efforts have been made to find a relationship contributing to this timing noise, such as sub-pulse drifts \\citep{2007A&A...469..607W}, but to no avail.  In this paper, using archival Arecibo observations, we have been able to reveal that this pulsar exhibits bursting episodes and confirm that it has moding behaviours. How this was performed and the implications are outlined in this paper in the following way. In Section 2 we provide technical information about the observations. In Section 3 we detail the data analysis techniques used to expose relationships from pulse-to-pulse. In Section 4 we conduct a fluctuation spectrum analysis on the normal and bursting modes found in Section 3. In Section 5, motivated by the recent X-ray observations of PSR~B0943+10, we make predictions for observations of B0611+22 at X-ray wavelengths that may help to understand the phenomenon seen here. Finally, in Section 6, we discuss our results and their scientific significance. ",
        "conclusions": "In summary, using a sensitive boxcar convolution method to preserve the phase resolution, we have found exceptional pulse patterns in the radio emissions from PSR~B0611+22 in archival data taken with the Arecibo radio telescope.  With this analysis, prolonged enhanced emissions were discovered at 327 MHz. When investigating the overall structure of this emission, it was shown that these are bursting events that abruptly appear and then systematically decay. These events appear to be superimposed upon a constant emission mode offset in pulse phase. This was unexpected and is a completely undocumented relationship for this pulsar.  To gain further insight into these events, we searched for them at 1.4 GHz using the same method, finding a different picture compared to 327~MHz. At 1.4~GHz, the bursting events were no longer bright episodes, and  the steady emission mode was no longer present. At this frequency PSR~B0611+22 appears to be a mode switching pulsar, where one mode is the decaying burst and the other mode is constant when on. This mode switching is surprising in this pulsar, only being suggested briefly in the past \\citep{1992msem.coll..280N,2000AAS...19713004N}. PSR B0611+22's characteristic age is also surprising, because the majority of mode-switching pulsars are a couple of orders of magnitude older \\citep{2007MNRAS.377.1383W}. This age discrepancy may suggest that bursting pulsars are a new classification  of pulsars that are independent of nulling and moding.  To investigate these relationships, we fitted Gaussians and measured the skewness of each convolved pulse. These measurements allowed us to confirm that there was a phase change between the two modes. This was used to isolate each of the mode statistics in each frequency range. We then showed, that the parameter differences were reflected in each mode's mean profile. These mean profiles, along with the skewness measurements, support that the modes are confined to the same phase region and therefore the same emission cone.  This analysis also shows that, apart from the amplitude difference from 327 to 1400 MHz, there was also a change in the distribution of the pulse-width parameters.  To see what could be causing this distribution change, we revisited the 327~MHz data and subtracted the constant underlying emission from the burst. When this was done, the normal mean profile from the 327~MHz bursting event distributions became consistent with the 1.4 GHz distributions.  An investigation of the pulse energy distributions also revealed differences between the two modes. While both modes can be described by log-normal probability density functions, at 327~MHz the modes are independent, but at 1400~MHz both mode energies are in a single function.  The two modes were also searched for drifting sub-pulses independently, where no significant signals were present in any of the fluctuation spectra, supporting \\cite{2006A&A...445..243W,2007A&A...469..607W} that sub-pulses are not prevalent in this pulsar. We have also shown that if B0611+22 acts similarly to other known pulsars, it could exhibit changes in its X-ray luminosity between the two separate modes. PSR~B0611+22 is particularly attractive as a potential X-ray source due to its relative youth. Correlated radio and X-ray observations are being planned for 2014.  It is interesting to note that PSR B0611+22 does exhibit period derivative variations in a manner analogous to those published by Lyne et al.~(2010; A.G.~Lyne, private communication based on Lovell radio telescope observations). It is currently an open question as to whether these changes are related to the bursting events we see here. Perhaps the pulsar bursts only in one of the spin-down rate changes? Our data, sampled on only a few days, are insufficient to answer this question. Further detailed studies are required.  Before this analysis, the curious emission properties of B0611$+$22 were not well known (Nowakowski 1992). Because this is the first pulsar that we have tried with this analysis, this may be a more common occurrence than we realise. In particular this work required the unique sensitivity of the Arecibo telescope to reveal this behaviour. Further studies of this kind are clearly necessary.  On a more global scale, studies such as this can help draw relationships between different types of modal and emission variations. Here it is easy to imagine a threshold situation that can relate nulling and bursting events. Consider, for example, similar pulsars to B0611$+$22 that are either not as bright or are farther away such that the normal mode is undetected. In this case the normal mode would be interpreted as being in a null state. A similar argument was made for PSR B0656+14 in the context of a possible connection with rotating radio transients \\cite{2006ApJ...645L.149W}. Further examples of this phenomenon will allow us to build up a more complete picture of the complex emission phenomenology of radio pulsars."
    },
    "1402/1402.4142_arXiv.txt": {
        "abstract": "The collapsar model explains the association of long duration gamma-Ray Bursts  (GRBs) with stellar collapse. It involves a relativistic jet that forms at the core of a collapsing massive star. The jet penetrates the stellar envelope and the prompt GRB emission is produced once the jet is well outside the star. Most current models for generation of relativistic jets involve Poynting flux dominated outflows.  We explore here the propagation of such a jet through a stellar envelope. The jet forms a bow shock around  it. Energy dissipation at the head of this shock supplies energy to a cocoon that surrounds the jet. This cocoon exerts pressure on the jet and collimates it. While this description resembles the propagation of a hydrodynamic jets there are significant qualitative differences. Two Strong shocks, the reverse shock that slows down the hydrodynamic jet and the collimation shock that collimates it, cannot  form within the  Poynting flux dominated jet. As a result this jet moves much faster and dissipates much less energy while it crosses the stellar envelope. We construct here a simple analytic model that explores, self consistently, the jet-cocoon interaction and dynamics. Using this  model we determine the properties of the jet, including its velocity, propagation time and shape. ",
        "introduction": "\\label{sec:introduction}  Gamma-ray bursts (GRBs) are short and intense bursts of low energy gamma rays. They involve powerful relativistic jets. % Long GRBs are associated with death of massive stars. The collapsar model\\footnote{Note that we use here a  general definition of the collapsar model in which it involves any central engine that launches a jet within a collapsing star. This is regardless of the specific nature of the central engine. } \\citep{1993ApJ...405..273W, 1999ApJ...524..262M} combined these two facts. According to this model a compact object is formed at the center of a star following its core collapse.  {This compact object launches a jet  that } drills a hole through the star and breaks out through the surface. The   observed gamma-rays are emitted due to some internal dissipation process  far from the surface of the star. Clearly the propagation phase of the jet within the star is an essential ingredient of the model and in recent years a lot of numerical \\citep[e.g.][]{2001ApJ...550..410M,2003ApJ...586..356Z,2007ApJ...665..569M, 2009ApJ...699.1261M,2013ApJ...777..162M} and analytic \\citep[e.g.][]{2003MNRAS.345..575M,2005ApJ...629..903L,2007ApJ...665..569M, 2011ApJ...740..100B} efforts was devoted to explore the propagation of a hydrodynamics jets within  stars.     {Most current relativistic jet  models are based in one way or another on a central  engine that generate a collimated Poynting flux dominated outflow. This is partially motivated due to analogy with AGNs. For AGNs, Poynting flux is the only available option. In GRBs a thermally driven jet (a fireball) is also thermodynamically possible. Still   it is generally expected that those will be less powerful than the electromagnetic ones \\citep[see e.g.][]{2013ApJ...766...31K}. It is therefore important to investigate the properties of a magnetic jet that propagates in the star, and whether its typical properties agrees with the observational constraints.  Collimated relativistic MHD jets in stellar environments were studied extensively both numerically and analytically under the approximation of an axi-symmetric, steady, non-dissipative flow \\citep[e.g.][]{2007MNRAS.380...51K,2009MNRAS.394.1182K, 2008MNRAS.388..551T,2008AIPC.1054...71T,2009ApJ...699.1789T, 2008ApJ...679..990Z, 2009ApJ...698.1570L,2010ApJ...725L.234L,2011PhRvE..83a6302L, 2010NewA...15..749T, 2012MNRAS.426..595K}. In these studies, however, either the shape of the jet or the profile of the confining medium was predetermined. In a case of a predetermined pressure \\citet{2009ApJ...698.1570L,2010ApJ...725L.234L,2011PhRvE..83a6302L} have shown analytically how the properties of the jet are determined by a given distribution of the confining pressure, assuming a steady state.  In realistic  outflows, however, the confining pressure is built during the course of the jet propagation, and its profile depends on the details of this propagation. Therefore the properties of the propagating jet and the confining pressure should be determined self consistently. Recently, \\citet{2013ApJ...764..148L} conducted an analytic, self consistent analysis for the propagation of a magnetic jet in a medium. Their model involves some constraining assumptions on the cocoon that result in a cylindrical jet, having a cross section radius of the order of the light cylinder radius of the central object. In additions they assume that kink instability grows in the jet on a time scale which is comparable to a few light crossing time of the jet width. This have lead them to conclude that the magnetic the jet will be disrupted by kink instability deep in the star, and transform to a hydrodynamic jet. However, as we show here, the jet is expected to be much wider and it is most likely that the kink instability has no time to develop in the jet before it breaks out of the star. Therefore, we assume here that the jet remains axi-symmetric.      We study,  the dynamics of a Poynting dominated jet that propagates inside the star. The dissipation of energy at the jet's head leads to the formation of a hot cocoon that surround the jet. The cocoon  applies pressure on the jet and collimates it. We build a self consistent analytic model that follows the time evolution of the jet and the cocoon. We show that the pressure in the cocoon is {typically} large enough to collimate the jet close to the source, so that different parts of the jet maintain strong causal connection with the jet's axis. In this case the poloidal magnetic field is comparable to the toroidal field in the proper frame of the jet. This leads to a  smooth transition of the jet material from a free expansion state, near the engine, to a collated state. We also show that the typical width of the propagating jet is of the order of a few 10s light cylinder radii, which imply that the  main body of the jet is stable to kink modes and that the jet is likely to survive  crossing  the star. We compare our results with recent numerical simulations of a highly magnetized jet \\citep{BT14} and show that there is a good agreement between the numerical and the analytic results. Finally, we discuss the differences between the propagation of a magnetic jet and a hydrodynamic jet,  focusing on the breakout time of the jet from the star. We show that the propagation velocity of magnetic jets is relativistic in most parts of the star (unlike hydrodynamic jets which propagate  typically at sub relativistic velocities).    The paper is structured as follow: In section 2 we describe the overall picture and lay down our basic assumptions. We calculate, In section 3,  the propagation velocity of the jet and show that it depends only on the properties of the confining medium near the jet's head. We then proceed to describe the geometry of the jet and it's cocoon (section 4), followed by a discussion on the conditions at the base of the jet and the collimation of the jet (section 5). In section 6  we discuss the stability of the jet and its survival inside the star. Finally, in section 7 we obtain  engine minimal activity time for a jet breakout. This quantity is of outmost importance as it might be related to the observed plateau in the long GRB duration distribution \\citep{2012ApJ...749..110B}. ",
        "conclusions": "\\label{sec:conclusions}  We have developed an analytic model for the propagation of a relativistic Poynting flux dominated jet in a dense medium and applied it to collapsar jets. {Namely, } to relativistic  jets produced within a collapsing star that   propagate in their massive star progenitor. {Note that we consider as collapsars all cases of GRB models  that involves jets that propagate though their massive stellar progenitor, regardless of the nature of the central engine that powers them.}    As the jet propagates it pushes the material in front of it, leading to the formation of a bow shock ahead of the jet. Matter that passes through this shock is heated and and forms a cocoon ,that applies pressure on the jet and collimates it. At the same time the cocoon propagates sidewards into the stellar envelope. Inside the star the cocoon pressure is large enough to collimate the jet close to the launching point so that the jet's outer boundary maintains a strong lateral causal contact with its center. In this case the jet is collimated smoothly without dissipation, and therefore almost all of the magnetic flux reaches the head.  We have shown that both the  size, $\\rhead$, and the proper velocity, $\\uh$   of the jet's head depend only on the ratio $(L_j/r_L^2) (1/\\rho_{ext})$. The first factor depends on the total luminosity  of the jet and the size of its central engine  while the latter depends only on the  external density profile.  These quantities do not depend on the details of the  jet structure behind its head.   Since the energy is injected into the cocoon only at the jet's head, this enables us to analytically estimate the energy injection rate into the cocoon without the need to know the  exact structure of the jet. Using this energy injection rate we  calculate  the cocoon's structure. Finally, once the cocoon pressure profile is known we determine the jet's structure. We show that the resulting jet structure is indeed consistent with the properties of the jet's head. In spite of the various approximations used, the analytic model is in a good agreement with the results (see Figure \\ref{fig:jet_simulation}) of recent  simulations \\citep{BT14}.   The propagation of the Poynting flux dominated jet described here should be contrasted with the propagation of a similar hydrodynamic jet. The basic difference arises due to the fact that while strong shocks can form in a hydrodynamic flow they cannot arise in a Poynting flux dominated outflow \\citep{1984ApJ...283..710K}. Therefore in a hydrodynamic jet the head of the jet involves both a forward shock that propagates into the surrounding stellar material and a reverse shock that propagates into the jet's material and slows it down significantly. At the same time the pressure exerted by the cocoon on the jet leads to a collimation shock that takes place deep near the base of the jet \\citep{1997MNRAS.288..833K,2007ApJ...671..678B,2011ApJ...740..100B}. This shock heats initially cold matter and the slows it down. The interplay between these two shocks, the one at the head of the jet and the collimation shock at the base determines the hydrodynamic jet structure. Unlike the Poynting flux dominated jet the hydrodynamic jet is slow and for most relevant parameters its head propagates with sub-relativistic velocities. This means that the jet crossing time is longer and the energy deposited by the shock and given to the stellar envelope is larger by an order of magnitude than the corresponding quantities for the corresponding Poynting flux dominated outflow.    The most questionable assumption is that the magnetic jet remains axi-symemtric and stable. Our estimates have shown that the kink instability, which is the fastest growing instability,  is unlikely to distrupt the jet while it propagates in the star. It could, however,  somewhat change the details of our results, for example  it may increase somewhat the size the head and therefore decreases its velocity and subsequently increases the jet escape time. However our analysis is based on linear estimates of  the growth rate of the instability and even those are somewhat approximate. Clearly only full three dimensional simulations could clarify the issue of stability.   \\citet{2013ApJ...764..148L}  have analyzed recently a similar system reaching a very different conclusions. They have argued that such a magnetic dominated jet becomes unstable while propagating within the star. Due to this instability the jet dissipates its magnetic energy on a time scale that is much shorter than the jet propagation time. The  jet continues propagating as a hydrodynamic jet within the star. The main difference between their analysis and ours  is that \\citet{2013ApJ...764..148L} assume  that the inner part of the cocoon is  dominated by the pressure of toroidal magnetic fields, exerting a much larger pressure on the jet boundary. In addition they assume that the cocoon pressure is uniform in the z direction. These two assumptions result in a cylindrical jet having a width that is comparable to $\\rhead$. Such a narrow jet is unstable to kink instability. Overall we expect that the assumption of a toroidal magnetic pressure dominated region in the cocoon is not entirely justified. In particular numerical simulations \\citep{BT14} do not show such a structure.        When applying these results to  collapsars we recall that the typical size of the head is  a few light cylinder radii $\\rL$  and the  propagation proper velocity is  mildly relativistic ($\\uh\\sim\\rhead/\\rL\\sim$ a few)  during most of the propagation inside the star. Thus, the jet crosses the star in a time that is close to its light crossing time, $R/c$. The fast head velocity implies that once the jet is launched it will most likely exit the star. The fast crossing time also implies that the energy injected into the star is minimal. Astrophysical implications of these finds will be discussed elsewhere \\citep{BGPL14}.    \\subsection*{Acknowledgements} We thank  A. Tchekhoskoy, M. Kuntz, A Spitkowsky, J. Stone and A. Levinson for helpful discussions and the  ISSI (Bern) for hospitality while the final version of this paper was written. This research was supported by the ERC advanced research grant ``GRBs'' by the  I-CORE (grant No 1829/12), by HUJ-USP grant and by the Max-Planck/Princeton Center for Plasma Physics.  \\appendix"
    },
    "1402/1402.6371_arXiv.txt": {
        "abstract": "We investigate the interacting dark sector composed of  dark matter, dark energy,  and dark radiation for a  spatially flat Friedmann-Robertson-Walker (FRW) background by introducing a three-dimensional internal space spanned by the interaction vector  $\\mathbf{Q}$  and solve the source equation for a linear transversal interaction. Then, we explore a realistic model with dark matter coupled to a scalar field  plus a decoupled radiation term, analyze the amount of dark energy in the radiation era and find that our model is consistent with the recent  measurements  of  cosmic microwave background anisotropy coming from Planck  along with   the future constraints achievable by CMBPol experiment. ",
        "introduction": "One could consider an exchange of energy between the components of the dark sector, so  dark matter not only can feel the presence of   dark energy through a gravitational expansion of the universe but also can interact between them \\cite{jefe1}, \\cite{T1}, \\cite{jefe4}, \\cite{ot}.  A coupling between dark energy and dark matter changes the background evolution of the dark sector allowing us to constrain any particular kind of interaction and giving  rise to a richer cosmological dynamics compared with non-interacting models \\cite{jefe1}, \\cite{NQ}. One way to extend the insight about the dark matter-dark energy interacting mechanism is to explore a bigger picture in which a third component is added, perhaps a  weakly interacting radiation term as it occurs within warm inflation paradigm. A step forward for constraining dark matter and dark energy  is to use  the physic behind  recombination or big-bang nucleosynthesis epochs by adding a decoupled radiation term to the dark sector for taking into account the stringent bounds  related to  the behavior of dark energy at early times \\cite{jefe4}, \\cite{hmi1}, \\cite{hmi2}. The authors have explored the behavior of a universe filled with  three interacting components, i.e., dark energy, dark matter, and dark radiation \\cite{jefe2} plus a decoupled radiation term (Model I). Then, they have extended the aforesaid scheme by considering a more realistic model with dark matter coupled to a scalar field plus a decoupled radiation (Model II) \\cite{jefe3}.  We perform a cosmological constraint using the updated Hubble data \\cite{H29},  bounds for  early dark energy \\cite{EDE1}, \\cite{Amendola},  \\cite{Cyburt}, \\cite{Cyburt2} and compare  our  constraints on cosmological parameters with the bounds reported by Planck  measurements \\cite{Planck2013} and WMAP-9 project \\cite{WMAP9} . ",
        "conclusions": "We will provide a qualitative estimation of the cosmological paramaters by constraining them with the Hubble data  \\cite{obs3}- \\cite{obs4} and the stricts bounds for the behavior of dark energy at early times \\cite{EDE1}-\\cite{EDE2}. In the former case, the  statistical analysis is based on the $\\chi^{2}$--function of the Hubble data \\cite{obs3}, \\cite{obs4}. For model I,  the Hubble expansion is given by \\be \\n{Ht} H_{\\rm I}(\\theta| z)=H_{0} \\Big( {\\cal A}x^{3\\ga_{s}}+{\\cal B}x^{4}+{\\cal C}x^{3}\\Big)^{\\frac{1}{2}}. \\ee Fron now on  we will  choose  $\\ga_{+}=4/3$ and  $\\ga_{-}=1$, so the parameters of the model are   $\\theta=\\{H_{0},\\ga_{s},  \\Omega_{x0},\\Omega_{m0}\\}$, and  will use that the flatness condition implies $\\Omega_{r0}=1-\\Omega_{x0}-\\Omega_{m0}$. \\begin{figure}[hbt!] \\begin{minipage}{1\\linewidth} \\resizebox{1.6in}{!}{\\includegraphics{2.eps}} \\resizebox{1.6in}{!}{\\includegraphics{3.eps}}\\hskip0.05cm \\resizebox{1.6in}{!}{\\includegraphics{4.eps}}\\hskip0.05cm \\resizebox{1.6in}{!}{\\includegraphics{5.eps}}\\hskip0.05cm \\caption{\\scriptsize{Model I: two-dimensional C.L. associated with $1\\sigma$,$2\\sigma$ for different $\\theta$ planes.}} \\label{Fig1} \\end{minipage} \\end{figure} We obtain that $\\ga_{s}$ varies from $10^{-4}$ to $10^{-3}$, so  these values clearly fulfill the constraint $\\ga_{s}<2/3$ that assure the existence of accelerated phase of the universe at late times [see Fig. (\\ref{Fig1})]. We find the best-fit at $(H_{0}, \\Omega_{x0})=(74.32 {\\rm km~s^{-1}\\,Mpc^{-1}},0.77)$ with $\\chi^{2}_{\\rm d.o.f}=0.783$ by using the priors $\\Omega_{m0}=0.2$ and $\\ga_{s}=10^{-3}$.  The value of $\\Omega_{x0}$ is slightly greater than the standard one of $0.7$ being such discrepancy only of $0.1\\%$. Moreover, we find that using the priors  $(H_{0}, \\ga_{s})=(74.20 {\\rm km~s^{-1}\\,Mpc^{-1}}, 10^{-3})$ the best-fit values for the present-day density parameters  are $(\\Omega_{x0},\\Omega_{m0})=(0.74, 0.23)$ along with a lower goodness condition ($\\chi^{2}_{\\rm d.o.f}=0.779$). In performing the statistical analysis, we find that $H_{0} \\in [71.28, 74.32]{\\rm km~s^{-1}\\,Mpc^{-1}}$ so the estimated values are met  within $1 \\sigma$ C.L.  reported by Riess \\emph{et al}, to wit,  $H_{0}=(72.2 \\pm 3.6){\\rm km~s^{-1}\\,Mpc^{-1}}$.  We also find that the $z_{t}$ (transition redshift) is of the order unity varying over the interval $[0.63, 0.75]$, such values are close to $z_{t}=0.69^{+0.20}_{-0.13}$ reported in \\cite{Zt1}-\\cite{Ztn}.  In the case of  model II, one finds that the Hubble expansion takes the following form: \\be \\n{Ht} H_{\\rm II}(\\theta| z)=H_{0} \\Big( {\\cal A}x^{3\\ga_{s}}+{\\cal B}x^{3\\ga_{+}}+{\\cal C}x^{3\\ga_{-}}+ {\\cal D}x^{3\\ga_{r}}\\Big)^{\\frac{1}{2}}. \\ee Here, we consider  $\\theta=\\{H_{0},\\ga_{\\rm s},  \\ga_{+}, \\ga_{-},\\ga_{\\rm r}, \\Omega_{\\rm  v0},\\Omega_{\\rm dr0}, \\Omega_{\\rm c0}\\}$ as the independent parameters to be constrained  for the model encoded in the Hubble function (\\ref{Ht}).  We start our statistical estimations by performing a global analysis on the eight parameters that characterize the model. In doing so, we find  that $\\chi^{2}(\\theta)$ reaches a minimum at  $\\{H_{0},\\ga_{\\rm s},\\ga_{+},\\ga_{-},\\ga_{\\rm r},\\Omega_{\\rm  v0},\\Omega_{\\rm dr0},\\Omega_{\\rm c0}\\}=(70.00, 0.010, 1.300, 1.009, 1.3400, 0.749, 0.00006, 0.199  )$ along with  $\\chi^{2}_{\\rm d.o.f}=20.172/(29-8)\\simeq 0.9605 <1$. We see that $\\ga_{s}$ varies from $ 0.010^{+0.007+0.055}_{-0.090-0.112}$ to $ 0.010^{+0.271+0.335}_{-0.189-0.303}$ at $1\\sigma-2\\sigma$ confidence levels, so  these values fulfill the constraint $\\ga_{s}<2/3$ at $95\\%$ C.L. that ensure the existence of accelerated phase of the universe at late times [ See Fig. (\\ref{Fig2})]. We find the best fit at $ (H_{0},\\Omega_{\\rm v0})= (70.00^{+4.99+6.07}_{-2.03-3.16} {\\rm km~s^{-1}\\,Mpc^{-1}}, 0.7299^{+0.057+0.0658}_{-0.0109-0.0238})$ with $\\chi^{2}_{\\rm d.o.f}=0.736$ by using the priors $(\\Omega_{\\rm c0}, \\Omega_{\\rm dr0},\\ga_{s},\\ga_{+},\\ga_{-},\\ga_{\\rm r})=(0.1999, 6\\times 10^{-5}, 0.010, 1.3000, 1.0099, 1.3400)$.  Moreover, we find that using the priors $( H_{0},\\Omega_{\\rm dr0},\\ga_{s},\\ga_{+},\\ga_{-},\\ga_{\\rm r})=(70.00 {\\rm km~s^{-1}\\,Mpc^{-1}}, 6\\times 10^{-5}, 0.010, 1.3000, 1.0099, 1.3400)$ the best-fit values for the present-day density parameters are considerably improved, namely, these turn give ($\\Omega_{\\rm v0}, \\Omega_{\\rm c 0})= (0.729^{+0.167+0.218}_{-0.102-0.152}, 0.239^{+0.210+0.355}_{-0.281-0.364})$ along with the same goodness condition ($\\chi^{2}_{\\rm d.o.f}=0.7368$) .  Notice that $\\Omega_{\\rm c0}$ varies from $0.239^{+0.086+0.101}_{-0.019-0.042}$ to $0.239^{+0.215+0.267}_{-0.091-0.115}$ at $68\\%$, $95\\%$ C.L., whereas $H_{0} \\in [70.00^{+2.57+3.54}_{-2.04-3.05}, 70.00^{+4.99+6.07}_{-2.03-3.16}]{\\rm km~s^{-1}\\,Mpc^{-1}}$; thus,  the latter values are consistent with the analysis of ACT and WMAP-7 data that gives $H_0 = 69.7 \\pm 2.5{\\rm km~s^{-1}\\,Mpc^{-1}}$. We find that the transition redshift is   $z_{\\rm t}=0.604^{+0.203+0.325}_{-0.025-0.043}$ at $1\\sigma, 2\\sigma$ C.L.,  close to ones reported in \\cite{Zt1}, \\cite{Ztn} or with the marginalized best-fit values $z_{\\rm t}=0.623^{+0.039 }_{- 0.052}$ listed in \\cite{mizt}, \\cite{mizt2}, \\cite{Zt2}.  \\begin{figure}[hbt!] \\begin{minipage}{1\\linewidth} \\resizebox{1.5in}{!}{\\includegraphics{c1.eps}}\\hskip0.05cm \\resizebox{1.5in}{!}{\\includegraphics{c2.eps}}\\hskip0.05cm \\resizebox{1.5in}{!}{\\includegraphics{c3.eps}}\\hskip0.05cm \\resizebox{1.5in}{!}{\\includegraphics{c4.eps}}\\hskip0.05cm \\caption{\\scriptsize{Model II: two-dimensional C.L. associated with $1\\sigma$,$2\\sigma$ for different $\\theta$ planes.}} \\label{Fig2} \\end{minipage} \\end{figure} The behavior of dark energy at early times can be considered as a new cosmological probe for testing dynamical dark energy models. The fraction of dark energy at recombination epoch should fulfill the bound $\\Omega_{\\rm ede}(z\\simeq 1100)<0.1$ in order to the dark energy model be consistent with the big-bang nucleosynthesis (BBN) data.   Some signals could arise from the early dark energy (EDE) models uncovering the nature of DE as well as their properties to high redshift, giving an invaluable clue to the physics behind the recent speed up of the universe \\cite{EDE1}. Then,  it was examined the current and future data  for constraining the amount  of EDE, the cosmological data analyzed has led to an upper bound of  $\\Omega_{\\rm ede}(z\\simeq 1100)<0.043$ with $95\\%$ confidence level (CL) in case of relativistic EDE while for a quintessence type of EDE has given $\\Omega_{\\rm ede}(z\\simeq 1100)<0.024$ although   the EDE component is not preferred, it is also not excluded from the current data \\cite{EDE1}. Another forecast for  the bounds of the EDE  are obtained with the  Planck  and CMBPol experiments\\cite{EDE2}, thus  assuming a $\\Omega_{\\rm ede}(a \\simeq 10^{-3}) \\simeq 0.03$  for studying the stability of this value, it found that $1\\sigma$ error coming from Planck experiment is $\\sigma^{\\rm Planck}_{\\rm ede} \\simeq 0.004$ whereas the CMBPol  improved this bound by a factor 4 \\cite{EDE2}.  For the model I, we find that at  early times the dark energy  changes rapidly with the redshift $z$ over the interval $[10^{3}, 10^{4}]$ , indeed,  around $z \\simeq 1100$ we find that $2.4 \\times 10^{-11} \\leq \\Omega_{\\rm x}  \\leq 2.6 \\times 10^{-10}$. However for the model II, we find that the dark energy   around $z \\simeq 1100$, $\\Omega_{\\rm \\phi} \\simeq 0.014$.  In  comparison with the SPT constraints on the early dark energy density over WMAP7 alone it was found that the $95\\%$ upper limit on $\\Omega_{\\rm ede}$ is reduced from $0.052$ for WMAP7-only to $0.013$ for WMAP7+SPT. This is a $38\\%$ improvement on the upper limit of $\\Omega_{\\rm ede} < 0.018$ reported for WMAP7+K11 \\cite{Reic}. The upper limit is essentially unchanged at $\\Omega_{\\rm ede} < 0.014$ for WMAP7+SPT+BAO+SNe. The $\\Omega_{\\rm ede} < 0.013$ bound from WMAP+SPT is the best published constraint from the CMB (see \\cite{Hou} and reference therein). Our findings point out that the model constructed here not only fulfill the severe bound of $\\Omega_{\\rm \\phi}(z\\simeq 1100)<0.018$ obtained from the measurements of CMB anisotropy from ACT and SPT \\cite{Reic}, \\cite{Amendola}, \\cite{Hou}  but also is consistent with the future constraints achievable by Planck and CMBPol experiments \\cite{EDE2} as well, corroborating that the value of the cosmological parameters selected before, through the statistical analysis made with Hubble data, are consistent with BBN constraints. Besides,  regarding the values reached by dark energy  around $z=10^{10}$ (BBN), we find that $\\Omega_{\\rm \\phi}=0.0025$ at $1\\sigma$level indicating the conventional BBN processes that occurred at temperature of $1 {\\rm  Mev}$ is not spoiled because the severe bound reported  for early dark energy $\\Omega_{\\rm \\phi}(z\\simeq 10^{10})<0.21$\\cite{Cyburt} or a strong upper limit $\\Omega_{\\rm \\phi}(z\\simeq 10^{10})<0.04$ \\cite{Cyburt2} are fulfilled at BBN.    \\begin{figure}[hbt!] \\begin{minipage}{1\\linewidth} \\resizebox{1.5in}{!}{\\includegraphics{H1.eps}} \\resizebox{1.5in}{!}{\\includegraphics{H2.eps}}\\hskip0.05cm \\resizebox{1.5in}{!}{\\includegraphics{H3.eps}}\\hskip0.05cm \\caption{\\scriptsize{Comparison of $H_{0}$, $\\Omega_{\\rm c0}$, and $\\Omega_{\\rm v0}$ measurements, with estimates of $\\pm 1\\sigma$ errors, from a number of different  methods. These are compared with the data provided by Planck and WMAP9.}} \\label{H1} \\end{minipage} \\end{figure}  We will perform a further  comparative analysis by taking into account several data sets along with  theirs  constraints on the most relavant parameters \\cite{Planck2013}, \\cite{WMAP9}, namely, the values  or bounds reported for  $H_{0}$, $\\Omega_{\\rm v0}$, $\\Omega_{\\rm c0}$, and $\\Omega_{\\rm ede}$.  The  Planck power spectra leads to a low value of the Hubble constant, which is tightly constrained by CMB data alone within the $\\Lambda$CDM model. From the Planck+WP+highL analysis, it found that $H_{0}   = (67.3\\pm 1.2) {\\rm km s^{-1} Mpc^{-1}}$ at $68\\%$ \\cite{Planck2013}. A low value of $H_{0}$   has been found in other CMB experiments, most notably from the recent WMAP-9 analysis. Fitting the base $\\Lambda$CDM model for the WMAP-9 data, it is found $ H_{0} =(70.0 \\pm  2.2) {\\rm km s^{-1} Mpc^{-1}} $ at $68\\%$ C.L. \\cite{WMAP9}. Then, our best estimation $ H_{0} =70.00^{+3.76}_{-2.02} {\\rm km s^{-1} Mpc^{-1}} $ at $68\\%$ C.L  is perfectly in agreement with the value reported by WMAP-9 project but  shows a slightly difference with the  Planck+WP+highL data, which is less than $0.04\\%$. In Fig. (\\ref{H1}), we show other observational estimations of the Hubble constant that includes the megamaser-based distance to NGC4258, Cepheid observations, a revised distance to the Large Magellanic Cloud, and so on (see \\cite{Planck2013}). For dark matter,  Wiggle-Z data give $\\Omega_{\\rm c0}=0.309^{+0.041}_{-0.035}$,  while Boss experiment increases  dark matter amount in $0.019\\%$; the joint statistical analysis with  6dF+SDSS+ BOSS+ Wiggle-Z data lead to $\\Omega_{\\rm c0}=0.307^{+0.010}_{-0.011}$ at $1\\sigma$ level \\cite{Planck2013}, showing a discrepancy with our  estimation of dark matter, $\\Omega_{\\rm c0}=0.239^{+0.086}_{-0.019}$,  no bigger than  $0.24\\%$  [see Fig. (\\ref{H1})].  For the fraction of dark energy,  Planck+WP data give  $0.685 ^{+0.018}_{-0.016}$ at $68\\%$ C.L,  whereas the joint statistical analysis on  Planck+WP+highL+ BAO  gives   $0.692\\pm 0.010$ at $1\\sigma$ level \\cite{Planck2013}. In our case,  we find that  $\\Omega_{\\rm v0}=0.7299^{+0.057}_{-0.0109}$, then the relative difference between Hubble data and Planck+WP+highL is $0.068\\%$ [see Fig. (\\ref{H1})].  Besides, the CMB anisotropies measurements put further  constraints on the behavior of dynamical dark energy in the recombination epoch; the latest constraints on early dark energy come from  Planck+WP+highL data and leads to $\\Omega_{\\rm ede} < 0.009$  at $95\\%$ C.L \\cite{Planck2013}. We have found that $\\Omega_{\\rm \\phi}(z\\simeq 10^{3})\\simeq 0.014$, so the relative error between both estimations is  $|\\Omega^{\\rm Us}_{\\rm \\phi}-\\Omega^{\\rm Planck}_{\\rm ede}|/|\\Omega^{\\rm Planck}_{\\rm ede}| \\simeq  0.5 \\%$. The latter disagreement can be reduced in a future by including additional measurements of Hubble data points or   other data sets \\cite{Xu}, allowing to improve the statistical analysis performed here."
    },
    "1402/1402.1767_arXiv.txt": {
        "abstract": "It has been recently pointed out that removing the axial symmetry in the  ``multi-angle effects'' associated with the neutrino-neutrino interactions for supernova (SN) neutrinos, a new multi-azimuthal-angle (MAA) instability would arise. In particular, for a flux ordering $F_{\\nu_e} > F_{\\bar\\nu_e} > F_{\\nu_x}$, as expected during the SN accretion phase, this instability  occurs  in the normal neutrino mass hierarchy. However, during this phase the ordinary matter density can be larger than the neutrino one, suppressing the self-induced conversions. At this regard, we investigate the matter suppression of the MAA effects, performing a linearized stability analysis of the neutrino equations of motion, in the presence of realistic SN density profiles. We compare these results with the numerical solution of the SN neutrino non-linear evolution equations. Assuming axially symmetric distributions of neutrino momenta we find that the large matter term strongly inhibits the MAA effects. In particular,  the hindrance becomes stronger including realistic forward-peaked neutrino angular distributions. As a result, in our model for a $10.8$~$M_{\\odot}$ iron-core SNe, MAA instability does not trigger any flavor conversion during the accretion phase. Instead, for a  $8.8$~$M_{\\odot}$  O-Ne-Mg core SN model, with lower matter density profile and less forward-peaked angular distributions, flavor conversions are possible also at early times. ",
        "introduction": "The  flavor evolution of supernova (SN) neutrinos, is strongly impacted by the self-induced effects, associated with instabilities induced by   the  the neutrino-neutrino interactions in the deepest stellar regions~\\cite{Pantaleone:1992eq,Qian:1994wh,Sawyer:2005jk,Duan:2005cp,Duan:2006an,Hannestad:2006nj,Fogli:2007bk,Fogli:2008pt} (see also~\\cite{Duan:2010bg} for a recent review). In this context a key ingredient  in the characterization of these  effects  is related to the  current-current nature of low-energy weak interactions, which implies a ``multi-angle term''~\\cite{Duan:2006an,Sawyer:2008zs} $(1-{\\bf v}_{\\bf p} \\cdot {\\bf v}_{\\bf q})$, where ${\\bf v}_{\\bf p}$ is the neutrino velocity~\\cite{Qian:1994wh}. Till recently,  all   studies  have assumed the axial symmetry for this term, that would then depend only on the zenith-angle. It has been shown that the ``multi-zenith-angle'' (MZA) term in some cases can hinder the maintenance    of the coherent oscillation behavior for different neutrino modes~\\cite{Duan:2006an,Raffelt:2007yz,Fogli:2007bk,EstebanPretel:2007ec,Sawyer:2008zs,Mirizzi:2010uz}.  A valuable tool to diagnose the possible instabilities of a dense neutrino gas is given by  the linearized stability analysis of the neutrino equations of motion. The linearized equations including generic azimuthal and zenith angle distributions for neutrinos were first worked out  in~\\cite{Sawyer:2008zs}. Then, the formalism for the azimuthal symmetric case was further developed in~\\cite{Banerjee:2011fj}. The stability method would  allow one to determine  the possible onset of the flavor conversions, seeking  for an exponentially growing solution of the eigenvalue problem, associated with the linearized equations of motion for the neutrino ensemble. Recently, this method has been applied in~\\cite{Raffelt:2013rqa} removing  the  assumption of axial symmetry in the $\\nu$ propagation. As a result, a   multi-azimuthal-angle (MAA) instability was found in addition to the  bimodal~\\cite{Samuel:1995ri} and MZA ones~\\cite{Raffelt:2013rqa}. In particular, it was considered a  neutrino ensemble with a strong excess of $\\nu_e$ over $\\bar\\nu_e$, as expected during the SN accretion phase (at post-bounce times $t_{\\rm pb} \\lesssim 0.5$~s). In this situation, the instability has been found in normal mass hierarchy (NH,  ${\\Delta m^2_{\\rm atm}} = m_3^2-m_{1,2}^2>0$), where the system would have been stable imposing   the  perfect axial symmetry. Subsequently, the role of this instability   has been clarified in~\\cite{Raffelt:2013isa,Duan:2013kba} with simple toy models. The discovery of the new MAA effects, has also motivated first numerical studies of the non-linear neutrino propagation equations in SN, introducing the azimuthal angle as angular variable in the multi-angle kernel,  in addition to the usual zenith angle~\\cite{Mirizzi:2013rla,Mirizzi:2013wda}. Remarkably it was considered  the $\\nu$ propagation  only along a radial direction, i.e.  a local solution  along a specific line of sight, under the assumption that the transverse variations of the global solution are small. In this approximation, for the  unstable case discussed above, MAA effects   lead in NH to spectral swaps and splits analogous  to what produced in inverted hierarchy (IH,  ${\\Delta m^2_{\\rm atm}}<0$) by the known bimodal instability~\\cite{Samuel:1995ri} also in a completely isotropic neutrino gas~\\cite{Raffelt:2007cb}.   All these recent works assume that the MAA effects can develop without any matter hindrance. Remarkably,  a crucial  ingredient  to be considered to determine the impact of self-induced flavor instabilities is the  ordinary matter term, associated with the net electron densities $n_e$ in SNe. As pointed out in~\\cite{EstebanPretel:2008ni} for the axial-symmetric case, when $n_e$ is not negligible with respect to the neutrino density $n_\\nu$, the large phase dispersion induced by the matter for $\\nu$'s traveling in different directions, would partially or totally suppress the self-induced oscillations through peculiar MZA effects. Recent studies of this case performed with realistic SN models, indicates that this situation is  realized during the supernova accretion phase (at post-bounce times $t_{\\rm pb} \\lesssim 0.5$~s). As a consequence, the self-induced flavor conversions found in IH in the axial-symmetric models are strongly inhibited~\\cite{Chakraborty:2011nf,Chakraborty:2011gd,Saviano:2012yh,Sarikas:2011am}.  From a preliminary schematic study done in~\\cite{Raffelt:2013rqa} it results that the matter density required to suppress the MMA instability in NH   is larger than the one necessary to suppress the self-induced conversions in IH for the axial symmetric case. Motivated by these previous results, we find it is  mandatory to understand  what is the role of the dense matter on the MAA effects during the accretion phase. We will use   as benchmark for the neutrino  and  matter density profiles the SN models from  recent long term simulations of core-collapse  explosions, performed by the Basel-Darmstadt model. These were already considered by some of us in~\\cite{Chakraborty:2011nf,Chakraborty:2011gd,Saviano:2012yh}. The plan of our work is as follows. In Sec.~2  we present the neutrino equations of flavor evolution  without imposing axial asymmetry. Then, we describe  the setup to perform the  stability analysis of the linearized equations of motion. In Sec.~3 we present the results of the stability analysis and we compare them with the numerical solution of the equations. We focus on the accretion phase for two different SN progenitor models. Finally, in Sec.~4 we comment on our results and we conclude. ",
        "conclusions": "We have performed  a dedicated study  of the matter suppression  of the MAA instability, connected with the axial symmetry breaking in the self-induced oscillations,  during the accretion phase   for two SN models with different progenitor masses. We characterize  the SN densities and the neutrino angular spectra with results from recent SN hydrodynamical simulations. We compared the linear stability analysis  with the numerical results of the flavor evolution  in which we have looked  at a local solution of the equations of motions along a specific line of sight. For the case of an iron-core 10.8 M$_{\\odot}$  we found that during the accretion phase the dominant matter density strongly suppresses the MAA instability. In particular, including realistic forward-peaked $\\nu$ angular distributions significantly reduces the strength of the $\\nu$-$\\nu$ interaction term. As a result, in this case the matter suppression of the self-induced flavor conversions would be complete. In the case of  a low-mass O-Ne-Mg SN with 8.8 M$_{\\odot}$ progenitor, where  the accretion phase is extremely short, the matter density profile is lower  and the $\\nu$ angular distributions less forward-peaked than in iron-core models. As a consequence, we found that also with realistic angular distributions flavor conversions   would be    possible at early times, producing in principle interesting time-dependent modulations.  Our analysis is complementary to  previous studies~\\cite{Chakraborty:2011nf,Chakraborty:2011gd,Saviano:2012yh}, where some of us explored the matter suppression of self-induced flavor conversions in inverted neutrino mass hierarchy, in axial-symmetric models (see also~\\cite{Sarikas:2011am}). The complete suppression of the self-induced effects in both the mass  hierarchies for iron-core SNe, implies that the neutrino signal during the accretion phase will be processed only by the ordinary Mikheyev-Smirnov-Wolfenstein   effect in the outer stellar layers. This effect would allow in principle  to distinguish the neutrino mass hierarchy through the study of the  rise time of the SN neutrino signal~\\cite{Serpico:2011ir}. The phenomenological importance of our findings motivates further studies with other SN models to confirm the generality of our results. In particular, an accurate characterization of the neutrino angular distributions seems necessary in order to get accurate predictions on the matter suppression. At this regard,  in~\\cite{Sawyer:2008zs} it has been shown that if  perfect cylindrical symmetry in the initial neutrino distributions were given up, then  super-fast flavor turnovers could occur. These effects cannot be tested within our spherically symmetric SN model. However, recently three-dimensional SN simulations have been carried on, characterizing the neutrino signal during the accretion phase. Surprisingly, a lepton-emission asymmetry among different flavor has been found~\\cite{Tamborra:2014aua}. In particular, the  electron (anti)neutrino fluxes show a dipole structure, while the $\\nu_x$ are almost spherically symmetric. We plan to investigate in a future work the role of the matter suppression in this flavor configuration, including also the not axisymmetric neutrino and matter angular distributions.  Self-induced flavor conversions associated with the MAA instability would still be possible for O-Ne-Mg SNe during the accretion phase, and possibly for iron core SNe  during the cooling phase, when the matter term becomes sub-dominant with respect to the neutrino-neutrino interaction term. In these situations, a self-consistent treatment of the neutrino equations of motion considering also the flavor evolution along the  transverse direction is still lacking. This would imply passing from an ordinary to a partial differential equation problem, adding a big layer of complication in the solution of the equations of motion. This effort is well motivated by the perspective of getting an accurate  characterization of the SN neutrino spectral features that would be observable in the planned large underground neutrino detectors~\\cite{Choubey:2010up}."
    },
    "1402/1402.4863_arXiv.txt": {
        "abstract": "The brightest ultraluminous X-ray source currently known, HLX--1, has been observed to undergo five outburst cycles.  The periodicity of these outbursts, and their high inferred maximum accretion rates of $\\sim{\\rm few}\\times 10^{-4}~M_\\odot~{\\rm yr}^{-1}$, naturally suggest Roche lobe overflow at the pericenter of an eccentric orbit.  It is, however, difficult for the Roche lobe overflow model to explain the apparent trend of decreasing decay times over the different outbursts while the integrated luminosity also drops.  Thus if the trend is real rather than simply being a reflection of the complex physics of accretion disks, a different scenario may be necessary.  We present a speculative model in which, within the last decade, a high-mass giant star had most of its envelope tidally stripped by the $\\sim 10^{4-5}~M_\\odot$ black hole in HLX--1, and the remaining core plus low-mass hydrogen envelope now feeds the hole with a strong wind.  This model can explain the short decay time of the disk, and could explain the fast decrease in decay time if the wind speed increases with time.  A key prediction of this model is that there will be excess line absorption due to the wind; our analysis does in fact find a flux deficit in the $\\sim 0.9-1.1$~keV range that is consistent with predictions, albeit at low significance.  If this idea is correct, we also expect that within tens of years the bound material from the original disruption will return and will make HLX--1 a persistently bright source. ",
        "introduction": "\\label{sec:intro}  The source 2XMM J011028.1--460421, commonly known as HLX--1, was identified by \\citet{2009Natur.460...73F} as being hosted by the z=0.0224 galaxy ESO 243--49.  At the $\\sim 93$~Mpc distance of this galaxy, the peak flux of this source corresponds to an isotropic bolometric luminosity of $\\sim 2\\times 10^{42}~{\\rm erg~s}^{-1}$.  The off-center location in the galaxy, the variability of the source, and its spectral properties make this the best current candidate for an intermediate-mass black hole (IMBH).  The mass of the hole has been estimated from spectral fitting, matching of spectral states as a function of Eddington ratio with stellar-mass sources, and a limitation of the luminosity to less than the Eddington luminosity, to be $M\\sim 10^{4-5}~M_\\odot$ \\citep{2011ApJ...734..111D,2011ApJ...743....6S,2012ApJ...752...34G,2012Sci...337..554W}.  The source has now been seen to undergo five outbursts with a period of very close to a year \\citep{2013MNRAS.428.1944S}.  Intriguingly, and as we discuss in Section~\\ref{sec:obs}, the fifth outburst was delayed by a couple of weeks compared to predictions based on the previously observed period.  Nonetheless, the most natural interpretation of the period is that it is an orbital period.  If so, the factor of $\\sim 40$ variation in flux over a cycle implies that the orbit is eccentric \\citep{2011ApJ...735...89L,2013MNRAS.428.1944S}.  The high implied peak mass accretion rate $\\sim {\\rm few}\\times 10^{-4}~M_\\odot~{\\rm yr}^{-1}$ is straightforwardly explained by Roche lobe overflow at the pericenter of the orbit.  If the system is undergoing Roche lobe overflow, then as we discuss in Section~\\ref{sec:Roche} the donor star must be a helium star, because the disk inflow time from the tidal radius for a main sequence star is too long, and from a white dwarf is too short, to explain the observed decay times of a few months.  In such a model, we would expect that the inflow time would remain approximately constant from outburst to outburst, with slightly longer observed decay times for outbursts that have smaller fluence because in the standard theory of geometrically thin disks the viscous time is larger for smaller accretion rates.  However, fits to the outbursts suggest that the decay times have decreased from $\\sim 6$ months to $\\sim 2-3$ months (see Section~\\ref{sec:obs}) and their fluences have dropped by a similar factor.  It could be that this is not a meaningful trend but is instead the result of complex interactions in the accretion disk.  If instead this is a real effect, then we must consider another scenario for HLX--1.  As an example of an alternate scenario, we explore a model in which the donor star transfers mass via a wind.  The advantages of the model are that it can naturally explain the time scale of decay and the rapid change in that time scale, and that because the donor star need not be close to the tidal radius the dynamical state of the system need not be fine-tuned.  The disadvantage of the wind model is that the required outflow rates of $\\sim 10^{-3}~M_\\odot~{\\rm yr}^{-1}$ are at least an order of magnitude larger than any known wind rate.  As a way around this, we suggest in Section~\\ref{sec:wind} that within the past several years a high-mass giant that was already close to its critical luminosity was tidally stripped by the IMBH.  This could push the remaining star to or above its new critical luminosity and lead to very high outflow rates.  The existence of such high-mass giants is plausible given the inferred young massive cluster that hosts HLX--1 \\citep{2012ApJ...747L..13F,2014MNRAS.437.1208F}.  In this scenario, the transient nature of the system is the key to the changing fluence and decay time.  In Section~\\ref{sec:obs} we present our data analysis of the outbursts. In Section~\\ref{sec:Roche} we discuss the Roche lobe overflow model and in particular the expected time scales for inflow and decay.  In Section~\\ref{sec:wind} we explore the wind accretion model and show that when the wind speed is comparable to the orbital speed at pericenter, the net angular momentum of the accreted matter can be low enough that the resulting disk is small.  Thus the decay time of the disk can be in the observed range, and a small increase in the wind speed can reduce the decay time by tens of percent. We also report on analysis of the data from this source that was motivated by the wind scenario: we expect significant extra narrow-line absorption in the system if the accretion is via a wind rather than through Roche lobe overflow.  Indeed, although high-quality X-ray spectra have shown no evidence for high neutral hydrogen columns, there is some evidence for narrow features at the expected amplitude, albeit at low significance, in the predicted $\\sim 0.9-1.1$~keV range.  Thus although there are too many uncertainties to claim clear confirmation of the prediction, this is an encouraging match with our scenario.  In Section~\\ref{sec:dynamics} we explore the possible role of stellar dynamics, and demonstrate that for reasonable stellar number densities the system should be unchanged during our few-year observational window.  We present our conclusions in Section~\\ref{sec:concl}, where we emphasize that continued monitoring with Swift will be needed to narrow down the possible explanations for this unique source.  In particular, we note that in our disrupted giant scenario we expect the bound material to return within tens of years, after which the source will be persistently bright for centuries. ",
        "conclusions": "\\label{sec:concl}  The magnitude and evolution of the decay time seen from the outbursts of HLX--1 are difficult to explain in a Roche lobe overflow model.  This motivates study of a wind accretion model, which as we have shown can naturally explain the observed decay times if the wind speed evolves slightly from outburst to outburst.  The price to pay for the wind model is steep: the required accretion rate is near the top end of what is seen from Wolf-Rayet stars, and possibly a factor of a few higher yet.  We suggest that the needed wind loss rate could have been provided by a recent tidal disruption of a massive red giant, which left behind a helium core with a small hydrogen envelope.  We also note that the wind model would imply that there is considerable extra absorption in the system; motivated by this prediction, we searched for and found tentative evidence of such absorption.  This provides encouraging but not definitive support for our model.  On the long term, we expect that the source will become persistently bright within roughly a decade to decades due to the accretion of returning bound matter.  Further monitoring of this source will be crucial to determine its nature."
    },
    "1402/1402.3996_arXiv.txt": {
        "abstract": "Ultra Steep Spectrum (USS) radio sources are one of the efficient tracers of High Redshift Radio Galaxies (HzRGs). To search for HzRGs candidates, we investigate properties of a large sample of faint USS sources derived from our deep 325 MHz GMRT observations combined with 1.4 GHz VLA data on the two subfields ({\\ie}VLA-VIMOS VLT Deep Survey (VVDS) and Subaru X-ray Deep Field (SXDF)) in the XMM-LSS field. The available redshift estimates show that majority of our USS sample sources are at higher redshifts with the median redshifts $\\sim$ 1.18 and $\\sim$ 1.57 in the VLA-VVDS and SXDF fields. In the VLA-VVDS field, $\\sim$ 20$\\%$ of USS sources lack the redshift estimates as well as the detection in the deep optical, IR surveys, and thus these sources may be considered as potential high-z candidates. The radio luminosity distributions suggest that a substantial fraction ($\\sim$ 40$\\%$) of our USS sample sources are radio-loud sources, distributed over redshifts $\\sim$ 0.5 to 4. ",
        "introduction": "\\label{s:intro} High-redshift radio galaxies (HzRGs) are found to be hosted in massive intensely star forming galaxies which contain large reservoir of dust and gas (Seymour et al. 2007). Also, HzRGs are often found to be associated with over-densities {\\ie}proto-clusters and clusters of galaxies at redshifts $\\sim$ 2 - 5 (Galametz et al. 2012). Therefore, identification and study of HzRGs allow us to better understand the formation and evolution of galaxies at higher redshifts and in dense environments. In the literature, the well known $z-{\\alpha}$ relation {\\ie}the positive correlation between cosmological redshift and the steepness of the radio spectrum has been exploited to search HzRGs (Ker et al. 2012). The low$-$frequency radio observations are found to be more advantageous in detecting faint Ultra Steep Spectrum (USS) radio sources (in turn HzRG candidates) as their flux density is relatively higher at low-frequency due to their steeper spectral index. Using our deep 325 MHz low-frequency GMRT observations (5$\\sigma$ $\\sim$ 800 $\\mu$Jy) in combination with the deep 1.4 GHz observations (5$\\sigma$ $\\sim$ 80 $-$ 100 $\\mu$Jy) over $\\sim$ 1.0 deg$^{-2}$ in the VLA $-$ VIMOS VLT Deep Survey (VVDS) field and $\\sim$ 0.8 deg$^{-2}$ in the Subaru X-ray Deep Field (SXDF) field, we derive a sample of USS sources and investigate their nature to find HzRGs candidates at submJy flux level. Furthermore, it is interesting to study faint USS sources down to submJy level, as the radio population at submJy level appears to be different than that at brighter end {\\ie}above few mJy (Smol{\\v c}i{\\'c} et al. 2008). ",
        "conclusions": "With a large sample of faint USS sources (${\\alpha}_{\\rm 325~MHz}^{\\rm 1.4~GHz}$ $\\leq$ -1) derived from our deep 325 MHz GMRT observations and 1.4 GHz VLA data, we have shown that the criterion of using USS source remains an efficient method to select high redshift sources, even at submJy flux densities. The available redshift estimates of our USS sources have medians of $\\sim$1.18 and $\\sim$ 1.57, in the VLA-VVDS and SXDF fields, respectively. A fraction ($\\sim$ 15$\\%$ $-$ 20$\\%$) of USS sources without redshift estimates also lack optical, IR detections, and can be potential high-redshift candidates. Our study shows that in addition to powerful HzRG candidates, faint USS population also constitutes relatively less radio powerful AGNs possibly hosted in obscured environments."
    },
    "1402/1402.3672_arXiv.txt": {
        "abstract": "We have built an electron-phonon coupling model to describe the behavior of the ${Nb_xSi_{1-x}}$ transition edge sensor (TES) bolometers, fabricated by electron-beam coevaporation and photolithography techniques on a 2-inch silicon wafer. The resistance versus temperature curves of several sensors with different thickness are measured with different bias currents, ranging from 200 nA to 10 ${\\mu}$ A, and the electron-phonon coupling coefficient and the electron-phonon thermal conductance are calculated herein. Our values are quite comparable with those in metallic TES samples of other groups using different measurement methods, while we are using the transition region of our TES sample to calculate the electron-phonon coupling interaction. ",
        "introduction": "Introduction}  We have developed bolometer matrices with transition edge sensors (TES) in NbSi. This novel design increases the sensitivity for each pixel, reduces the phonon noise below ${10^{-17}W/\\sqrt{Hz}}$ \\cite{fabrication} and allows multiplexed readout by SQUIDs \\cite{multiplexer,book:SQUID}. The sensitive bolometer matrices and the transition edge sensors (TES) could be applied to measure the temperature anisotropies of the cosmic microwave background (CMB) radiation, which was created about 400,000 years after the Big Bang when the universe became transparent. The last-scattered photons of CMB radiation by free electrons become polarized by Thompson scattering. By study of the CMB temperature and polarization anisotropies, the initial condition of our Universe can be revealed and cosmological theories such as inflation or gravitational waves could be tested.  Our prototype 23 pixel ${Nb_xSi_{1-x}}$ bolometer matrices (with x=0.14) are fabricated on a 2-inch silicon wafers using electron-beam co-evaporation and photolithography techniques. The superconducting transition temperature can be adjusted by the Nb concentration \u201cx\u201d and the thickness of the ${Nb_xSi_{1-x}}$ thin film. Several sensors of the NbSi TES matrices are measured and observed to undergo metal to superconductor transition at 75 mK and 110 mK for the thickness of 20 nm and 50 nm, respectively. The resistance versus temperature curves, slightly shift to lower temperature values when a higher excitation current is applied.  We use an ``electron-phonon coupling'' model to describe our TES\\textquoteright s behavior. The electron-phonon coupling coefficient, ${g_{e-ph}}$, is calculated from the slope of the fitted curve of the electrically dissipated power into the TES versus (${T_e^5-T_{ph}^5}$), where ${T_e}$ is the electron temperature and ${T_{ph}}$ is the phonon temperature of the NbSi.  We compare our data results with other types of thermometers. ",
        "conclusions": ""
    },
    "1402/1402.1677_arXiv.txt": {
        "abstract": "We examine the regime of validity of $N$-point spectra predictions of single field inflation models that invoke transient periods of non-adiabatic evolution.  Such models generate oscillatory features in these spectra spanning frequencies up to the inverse time scale of the transient feature.  To avoid strong coupling of fluctuations in these theories  this scale must be at least $\\sim 10^{-2}/c_s$ of the Hubble time during inflation, where $c_s$ is the inflaton sound speed.   We show that, in such models, the signal-to-noise ratio of the bispectrum is bounded from above by that of the power spectrum, implying that searches for features due to non-adiabatic evolution are best focussed first on the latter. ",
        "introduction": "\\label{sec:intro}  In this work we examine the regime of validity of models of single field inflation that invoke transient periods of non-adiabatic evolution to generate features in the spectra of curvature fluctuations. We clarify the limits that can be placed on the width of a feature in the inflationary potential or sound speed, and more generally, the shortest time scale or highest energy scale that can be probed by inflationary fluctuations while the effective theory remains weakly coupled. The characterization of the region of parameter space within which perturbation theory is under control has become more important since the release of the Planck data. These data now accurately probes a range of scales large enough to encompass modes that were on the horizon and those that were above the strong-coupling scale at a single epoch.  Thus, the physical interpretation within the inflationary context of correlations between these scales must be treated with care.   Features in various aspects of the inflationary Lagrangian which generate a period of non-adiabatic evolution of the inflationary background are by now well established as a phenomenological way of introducing oscillatory signatures into the spectra of curvature fluctuations during inflation.  Starobinsky first noted that a potential containing a singularity in its first derivative lead to oscillations in the resulting spectrum of  curvature fluctuations \\cite{Starobinsky:1992ts}. More recently Starobinsky's model has been extensively revisited \\cite{Arroja:2011yu, Martin:2011sn,  Arroja:2012ae}. A step in the potential of canonical inflation results similarly in oscillations of the power spectrum and bispectrum as shown both numerically \\cite{Adams:2001vc, Chen:2006xjb, Chen:2008wn} and  analytically \\cite{Dvorkin:2009ne, Adshead:2011bw, Adshead:2011jq, Adshead:2012xz, Miranda:2012rm, Adshead:2013zfa, Miranda:2013wxa}. Oscillations in the potential lead to resonant logarithmic oscillations in these spectra \\cite{Chen:2008wn} and occur naturally in axion monodromy inflation \\cite{Silverstein:2003hf,Flauger:2009ab}. Integrating out an orthogonal field in a multi-field inflationary model with a turning trajectory generates oscillatory signals \\cite{Achucarro:2012sm, Achucarro:2012yr, Achucarro:2012fd, Achucarro:2013cva}. Further work on  the effects of non-adiabatic evolution during slow roll inflation includes \\cite{Chen:2010bka, Chen:2011tu, Chen:2011zf}.  Observationally, features were first invoked to explain observed glitches in the angular spectrum of temperature fluctuations in the cosmic microwave background (CMB) in the year one analysis of the Wilkinson Microwave Anisotropy Probe (WMAP) data \\cite{Peiris:2003ff}. Further studies were carried out fitting for these glitches \\cite{Covi:2006ci, Hamann:2007pa}. A hint of evidence for a high frequency oscillations due to a step feature in the inflationary potential was found in WMAP year 7 data by \\cite{Adshead:2011jq} and at a similar significance in the Planck temperature power spectrum \\cite{Ade:2013rta} with the same amplitude \\cite{Miranda:2013wxa}. Evidence for oscillations due to a resonant enhancement of the fluctuations has also been reported in WMAP data by \\cite{Peiris:2013opa}, however, the evidence weakened in the Planck data \\cite{Ade:2013rta,Meerburg:2013dla,Easther:2013kla}. The Planck polarization power spectrum should soon provide a more definitive empirical check of these interpretations \\cite{Mortonson:2009qv}.  In this paper, we examine theoretical limits both on these interpretations and on the effects of non-adiabatic evolution in general by requiring that perturbation theory remains valid throughout all aspects of their calculation. We detail how strong coupling of the fluctuations ultimately limit the shortest time scale associated with these phenomena and hence the parameters of any given model for inflaton features.  These results  parallel those found in the context of slow roll inflation \\cite{Cheung:2007st, Cheung:2007sv, Leblond:2008gg, Shandera:2008ai, Baumann:2011dt, Barnaby:2011pe}, here generalized for transient violations of slow-roll.  This work is organized as follows. In \\S \\ref{sec:perttheory} we study the breakdown of perturbation theory during inflation by considering the non-linearities of the equations of motion for the perturbations. In \\S \\ref{sec:action} we relate this breakdown to the emergence of strong coupling from the perspective of the action for the fluctuations. In \\S \\ref{sec:bispecandsnr} we examine bounds on the maximal bispectrum consistent with perturbation theory, and its resulting signal to noise. In \\S \\ref{sec:examples} we collect some examples of non-adiabatic evolution during inflation and its impact on inflationary fluctuations. Finally we conclude in \\S \\ref{sec:conclusions}.  Throughout, we work in natural units where the reduced Planck mass $M_{\\rm Pl} = (8\\pi G_N)^{-1/2} = 1$ as well as $c = \\hbar =  1$. ",
        "conclusions": "\\label{sec:conclusions}  In this work we have considered the bounds on non-adiabatic evolution of inflaton fluctuations during inflation. We have identified four frequency scales below which perturbations must lie in order to remain perturbative: the symmetry breaking scale, the $c_s<1$ strong coupling scale, the excitation energy conservation scale, and the excitation strong coupling scale. The first two scales place a strict lower bound on the width of any feature in the inflationary history for which observables can be reliably computed. This bound is independent of the amplitude of the feature, and represents the highest energy scale, or equivalently, shortest time scale at which the theory of the fluctuations is linear. Specifically, this time scale is approximately $10^{-2}/c_s$ of the Hubble time during inflation. While, it might appear that sufficiently small perturbations should always be under control, perturbation theory itself has broken down beyond this scale, and observables cannot be reliably computed.  For example, a statistical description of the curvature fluctuations in terms of $N$-point functions may fail to form a perturbative series.  Although the three-point function or bispectrum may be bounded for a small amplitude feature in the potential, higher order operators imply an increasing amount of non-Gaussianity as $N$ increases, with $\\mathcal{L}_N/\\mathcal{L}_2$ becoming order unity for a feature width at the strong coupling scale.   In a low sound speed model $\\mathcal{L}_3/\\mathcal{L}_2$ itself exceeds unity as modes pass the $c_s<1$ strong coupling scale even without a feature. Without a feature, of course, the appearance of such a scale is not a concern, provided there remains a hierarchy between this scale and the Hubble scale where all correlation functions are determined. One can then invoke the decoupling theorem to infer independence of the low energy observables from the physics above this scale. The difference in the situation at hand is that the non-adiabatic evolution may be thought of as exciting particle content at an energy scale set by the width of a feature. If this scale is above the strong coupling scale, one is in violation of the assumptions of the decoupling theorem since modes are not entering the theory in their adiabatic vacuum state.  For large amplitude features, where fractional deviations in the power spectrum approach unity, the excitations themselves violate perturbativity due to their amplitude and set the remaining two scales, the excitation energy conservation scale and the excitation strong coupling scale.  These determine the maximal non-Gaussianity allowed, bounding the maximal SNR in the bispectrum.  Nonetheless for the detection of a non-adiabatic feature, the {\\em{relative}} SNR in the bispectrum can at most be comparable to the power spectrum if the calculations are to remain perturbative.  While the above considerations may seem mainly of academic interest,  current CMB data now have the dynamic range  to encompass both superhorizon and strongly-coupled modes at a single epoch. Intriguingly these observations allow and even mildly favor correlated features extending to the strong coupling scale. Convincing observational evidence for apparent features that saturate or exceed these strong coupling bounds  may point to new physics, or at least call into question the self-consistency of the scenarios that predict them."
    },
    "1402/1402.7321_arXiv.txt": {
        "abstract": "Measurements of the total logarithmic central slope of the mass profile in galaxy clusters constrain their evolution and assembly history and that of their brightest cluster galaxies. We report the first full surface brightness distribution modelling of the inner region of the galaxy cluster MACS J1149.5+2223. We compare these results with a position-based modelling approach for which we employ more than twice the previously known positional constraints. This is the first time that the detailed lensed image configuration of two non-central cluster galaxies with Einstein rings has been mapped. Due to the extended radial coverage provided by the multiple images in this system, we are able to determine the slope $\\partial \\log{ \\kappa  }/\\partial \\log{R} = -0.33$ of the total projected mass distribution from $8$ to $80~\\mathrm{kpc}$. This is within the cluster-to-cluster scatter estimates from previous cluster measurements. Our reconstruction of the image surface brightness distribution of the large central spiral galaxy has a root mean square residual for all image pixels of $1.14~\\sigma$, where $\\sigma$ is the observational background noise. This corresponds to a reconstruction of the positions of bright clumps in the central galaxy with an rms of $0.063~\\mathrm{arcsec}$. ",
        "introduction": "Our current understanding of structure formation is based on a hierarchical picture where more massive dark matter (DM) structures form via the infall and the progressive merging of smaller objects. DM-only simulations, within the cold dark matter (CDM) paradigm, have predicted, over the last two decades, remarkably self-similar DM mass profiles that are well fitted by gradually steepening analytic models, such as the Einasto \\citep{Einasto_Profile} or the Navarro, Frenk and White \\citep[NFW;][]{NFW_Profile} profile, over a wide range of scales from galaxy clusters \\citep{Springel_cluster_2001, Diemand_cluster_2004, Reed_cluster_2005,Springel_nature_Millenium_2005,Merritt_2006,Gao_Phoenix} to galaxies \\citep{Navarro_Aquarius_galaxy_2010}. A central slope $\\gamma_1 = \\partial \\log{\\rho} / \\partial \\log{r} \\approx  -1$ and a steeper outer slope $\\gamma_2 \\approx -3$ are generally found in those simulations. However, the details of the mass distribution are still being refined, and in general, there is no definitive prediction from numerical simulations on the asymptotic central slope of the galaxy cluster mass distribution. In a suite of nine high-resolution simulated N-body DM-only clusters, \\cite{Gao_Phoenix} have found considerable ($\\sim 20$ percent) halo-to-halo scatter. On scales of a few $\\mathrm{kpc}$, the baryonic component of the brightest cluster galaxy (BCG) and the dark matter halo play comparable roles in shaping the total density. However, there are competing hypotheses about their relative importance and the relevant physical processes involved. The inner DM profiles can be steepened by adiabatic contraction \\citep{Gnedin_2004, Sellwood_2005, Sommer-Larsen_2010}, or flattened by (repeated) ejection of gas by active galactic nuclei (AGN) \\citep{Martizzi_2013} or by heating by dynamical friction \\citep{ElZant_2001,ElZant_2004}. Even in simulations with dissipationless mergers of multi-component systems (DM+stars), a flattening of the inner DM profile can be achieved \\citep{Laporte_2012}.  Gravitational lensing provides a powerful tool to compare the predictions of cluster simulations with observations. There has been a considerable effort to get a coherent picture of the total cluster mass distribution on a large range of scales by combining stellar kinematics at the innermost radii ($r<10~\\mathrm{kpc}$) with weak lensing measurements at radii $r>100~\\mathrm{kpc}$ and strong lensing on intermediate scales \\citep[e.g.][]{Sand_2004,Sand_2008,Newman_2013a}. In particular, \\cite{Sand_2008} combined stellar velocity dispersion observations with strong-lensing data for the cluster Abell 383 and found that the inner slope of the dark matter is $\\gamma_{\\mathrm{DM}} \\approx -0.45$. This is shallower than the values found at similar radii in purely DM simulations. More recent observational constraints come from the analysis by \\citet{Newman_2013a}. These authors have performed a combined stellar kinematic, strong lensing and weak lensing mass reconstruction for seven massive $(0.4 < M_{200}/ (10^{15} \\Msun) <2)$\\footnote{$M_{200}$ is the mass within a sphere with radius $r_{200}$ and a mean enclosed density 200 times the critical value} clusters with redshifts between 0.2 and 0.3. They find a 3D central logarithmic total mass density slope of $\\gamma \\approx -1.16$ consistent with the NFW profile and with a scatter between clusters of $\\sigma_\\gamma = 0.13 \\left( 68 \\% \\mathrm{CL}\\right)$. This, however, is the total density profile, and the central $<10~\\mathrm{kpc}$ are typically dominated by the stars of the BCG. An analysis of the separate contributions of baryons and DM in \\citet{Newman_2013b} indicates a flattening of the DM profile with a logarithmic central slope of $\\gamma_{\\mathrm{DM}} \\approx -0.5$.  In this paper, we focus on building a high resolution lens model for \\macswo~and on comparing the slope inferred from the lensing reconstruction with the predictions for the total central mass density slope of galaxy clusters from cosmological structure formation simulations.  The cluster \\macswo~was initially discovered as part of the Massive Cluster Survey (MACS) quest for the most massive galaxy clusters in the Universe \\citep{Ebeling01,Ebeling07}. More recently, \\citet{Zheng_2012_magnified} have used the magnifying power of this cluster to identify a faint galaxy located at redshift $z=9.6\\pm0.2$.  So far, two gravitational lens models for the cluster \\macs have been published. The first model, by \\cite{Zitrin_2009}, assumes that the mass approximately follows the light. The model consists of a superposition of power-law mass profiles for each galaxy in the cluster. As constraints, this model uses the multiple image positions from strong lensing. However, many of the details, especially for system 1 (see Sec.~\\ref{sec:observations_and_constraints}), are only reproduced approximately. Under the assumption that the BCG consists only of stellar mass, this models infers a nearly uniform DM surface mass density out to $\\sim 200~\\mathrm{kpc}$.  The second model, by \\cite{Smith_2009}, is also based on the image positions of multiply lensed bright clumps. However,  it makes us of a larger number of bright clumps, hence a significantly larger number of constraints. The reconstructed image positions have a root mean square (rms) deviation from the observed positions of $0.5 \\asec$. This model rules out the flat central profile proposed by \\cite{Zitrin_2009}.  In this work, we improve on these previous models in three aspects. First, we use a more sophisticated model. We model {\\em all five} galaxies that are close to multiply lensed images and close to the cluster centre using individual mass profiles. This is crucial for reproducing the morphology of the lensed image of the main system (see Sec.~\\ref{sec:observations_and_constraints}). Secondly, we identify twice as many positional constraints as previously used. These include multiply-imaged clumps that are part of two Einstein rings formed by two cluster galaxies, as well as details of the nonlinear configuration of the image covering the centre of the cluster. Thanks to the increased number of constraints, we can place tight constraints on the slope of the total mass density profile. Thirdly, we use a more sophisticated gravitational lens modelling method that was originally developed and applied on galaxy scales by \\citet{Koopmans_2005, Suyu_2006,Vegetti_2009} and recently also on cluster scales by \\cite{Eichner_2013_macsj1206}. In this method, we use the information provided by the positions of multiply-lensed images along-side with the full surface brightness distribution of the images.  This paper is organized as follows. We describe the observations, the image morphology of the main lensed image and the newly identified positional constraints in section \\ref{sec:observations_and_constraints}. In section \\ref{sec:mass_model} we describe our analytic mass parametrization and in section \\ref{sec:modelling_methods} our two modelling methods. Section \\ref{sec:results} contains the main results, the best lens models for the position modelling and the surface brightness modelling in sections \\ref{sec:results_modelling_pos} and \\ref{sec:results_modelling_brightness} respectively. The slope measurement of the total mass distribution is presented in section \\ref{sec:Central Slope}. Throughout this paper, we assume $H_0 = 67.3$ and $\\Omega_{\\mathrm{m}} = 0.315$ from Planck \\citep{Planck_Cosmology}. At the redshift of the cluster, $z = 0.544$, 1 arcsec corresponds to $~6.6$~kpc. ",
        "conclusions": "We have presented a new and detailed model for the centre of the galaxy cluster \\macswo. In particular, we have identified more than twice as many constraints as previously used.  We have also used a multiple lens plane algorithm  in order to properly include the lensing contribution of the mass associated with the source S1. Finally, with a hybrid modelling approach, we have performed the first detailed reconstruction of the surface brightness distribution of the system 1. Our results can be summarized as follows:  \\begin{enumerate} \\item We have recovered the surface brightness distribution of system 1 with a precision close to the noise level of the HST CLASH observations. \\item Thanks to the hybrid modelling approach, we have derived posterior probability density distributions of the main model parameters that are significantly tighter than those derived with the simpler position modelling. \\item Thanks to the new constraints, we have constrained three important details of the mass distribution: the individual mass distributions of the two cluster galaxies G1 and G2, and the total mass distribution of the cluster at the innermost radii. \\item  We  have recovered the 2D logarithmic slopes $\\gamma_{\\mathrm{2D}} = \\partial \\log \\left(\\kappa \\right)/ \\partial \\log  R  \\approx -2.51$ and $ \\approx -1.63$ for galaxies G1 and G2 measured at a distance corresponding to their respective Einstein radius in the cluster. \\item Our mass model suggests a large $(\\sim 12~\\mathrm{arcsec})$ core in the cluster DM distribution and that the total mass profile at the very centre of the cluster is dominated by the BCG. We have found a central logarithmic slope of the 2D mass distribution between $1.2$ and $12~\\mathrm{arcsec}$ of $\\gamma_{\\mathrm{2D}} = \\partial \\log \\left(\\kappa \\right)/ \\partial \\log  R  \\approx -0.39$ for the DM halo+BCG mass components and  $ \\approx -0.33$ when the other central galaxies and cluster members are included.  \\end{enumerate}"
    },
    "1402/1402.2218_arXiv.txt": {
        "abstract": "{ We present an extensive analysis on the determination of the isotropic radio background. We consider six different radio maps, ranging from 22 MHz to 2.3 GHz and covering a large fraction of the sky. The large scale emission is modeled as a linear combination of an isotropic component plus the Galactic synchrotron radiation and thermal bremsstrahlung. Point-like and extended sources are either masked or accounted for by means of a template. We find a robust estimate of the isotropic radio background, with limited scatter among different Galactic models. The level of the isotropic background lies significantly above the contribution obtained by integrating the number counts of observed extragalactic sources. Since the isotropic component dominates at high latitudes, thus making the profile of the total emission flat, a Galactic origin for such excess appears unlikely. We conclude that, unless a systematic offset is present in the maps, and provided that our current understanding of the Galactic synchrotron emission is reasonable, extragalactic sources well below the current experimental threshold seem to account for the majority of the brightness of the extragalactic radio sky. }  \\date{\\today} \\begin{document} ",
        "introduction": "\\label{sec:Intro} Two main avenues can be followed to estimate the total brightness of the extragalactic radio sky. The first and most obvious way is to simply collect all the radio waves arriving from the sky in a detector, which, nowadays, typically is a single dish telescope or a balloon/satellite. The main problem of this estimate resides in the subtraction of the foreground emission from within our Galaxy. The second way consists instead in adding up the single contributions from all the extragalactic sources. It can be done by individually observing the sources or by statistically evaluating their contribution from the fluctuations of the intensity below the detection threshold. To this aim, synthesis arrays are the most powerful telescope currently developed. Their interferometric working mechanism prevents, on the other hand, the detection of a truly isotropic emission.  By themselves, both methods cannot provide an ultimate answer to the radio background estimate. In the first case, a degeneracy with other contributions, like the monopole of the cosmic-ray (CR) emission in the Milky-Way, a zero-level offset in the maps, or local emissions, is unavoidable. In the second case, since the experimental sensitivity is finite, there is always an open window for a contribution from faint populations of sources. Our understanding of the total brightness of the extragalactic radio sky can be settled down, only if the estimates from the two techniques match. Unfortunately, at present, they do not match.  Indeed, the balloon--borne experiment ARCADE 2 (Absolute Radiometer for Cosmology, Astrophysics and Diffuse Emission)~\\cite{Singal:2009xq} recently reported an isotropic radio emission (based on new observations at 3 to 10 GHz and on a re-analysis of old maps at lower frequencies) which significantly exceeds the expected contributions from known extragalactic sources. The isotropic background has been estimated by subtracting a Galactic component. The latter was computed employing two different methods: using a simple Galactic plane-parallel model with a $cosec(b)$ dependence on the galactic latitude; correlating the radio emission with a map of the CII recombination line. These two techniques give consistent results.  The goal of the paper is to revisit the isotropic background estimate, employing a more sophisticated analysis, to assess the robustness of such excess. In particular, we will use state-of-the-art models for the Galactic foreground.  The source contributions to counts and to the sky brightness are linked by the Rayleigh-Jeans law: \\be \\frac{dN}{dS}(S)\\,S\\,dS=\\frac{2\\,k_B}{c^2}\\,\\nu_{\\rm obs}^2\\,dT\\;, \\ee where $dN/dS$ is the number of sources per steradian and per unit flux in the range of flux $(S,S+dS)$, $\\nu_{\\rm obs}$ is the radio-frequency of observation, $T$ is the brightness temperature, $k_B$ is the Boltzmann constant, and $c$ is the speed of light. The total brightness of the extragalactic sky is thus given by: \\be T_E=\\int\\ dT=\\frac{c^2}{2\\,k_B\\,\\nu_{obs}^2}\\ \\int_{0}^{S_{\\rm max}}\\ dS \\ \\frac{dN}{dS}(S)\\,S\\;, \\label{eq:totT} \\ee where $S_{\\rm max}$ is the flux density of the brightest sources. However, observationally we can obviously measure source counts only down to a certain threshold $S_{\\rm thr}>0$ because of finite sensitivities. Thus, part of the integral in Eq.~\\ref{eq:totT} has to rely on an extrapolation of $dN/dS$ at low fluxes.  The most recent estimate of $T_E$ from number counts is presented in Ref. \\cite{Vernstrom:2013vva} from 1.4 GHz data, with the $dN/dS$ derived down to the $\\mu$Jy level. The extrapolation to lower fluxes does not affect the isotropic temperature determination if such extrapolation is assumed to follow from standard evolutionary models, like the one in Ref. \\cite{Condon:1984}. The reported background temperature in Ref. \\cite{Vernstrom:2013vva} is 115 mK.   On the other hand, by subtracting a model for the foreground Galactic emission~\\cite{Fixsen:2009xn}, the ARCADE Collaboration isolated, in radio maps from 22 MHz to 10 GHz, an isotropic temperature (on top of the CMB blackbody contribution), which can be fitted by a power law $T_E=T_s \\left(\\nu/{\\rm GHz}\\right)^{\\alpha}$ with $\\alpha=-2.62\\pm0.04$ and $T_s=1.19\\pm 0.14$ K~\\cite{Fixsen:2009xn}. This leads to about 500 mK at 1.4 GHz, so exceeding the counts-based estimate by a factor of 4-5.  Various astrophysical solutions for the mismatch have been discussed. The contribution from radio supernovae, radio quiet quasars and diffuse emission from intergalactic medium and clusters have been shown to be quite modest \\cite{Singal:2009dv} and cannot provide a sufficiently high bump in the sub-$\\mu$Jy range of number counts. Star forming galaxies with a non-standard evolutionary model are constrained by bounds from the far IR-radio correlation, gamma-rays, and P(D) analyses \\cite{Singal:2009dv,Lacki:2010uz,Vernstrom:2011xt,Ponente:2011se,Condon:2012ug,Vernstrom:2013vva}. The required level of contribution can be also efficiently probed by anisotropy studies in the case of clustered sources~\\cite{Holder:2012nm}. If the origin of the excess is from a population of extragalactic sources, they have to be very faint (significantly contributing to number counts well below the $\\mu$Jy level) and very numerous \\cite{Singal:2009dv,Vernstrom:2011xt,Vernstrom:2013vva}.  A possibility in this direction is represented by synchrotron radiation induced by annihilations of dark-matter particles in halos of extragalactic structures, as proposed in Ref. \\cite{paper1}. This has been discussed also in Refs. \\cite{Fornengo:2011xk,Hooper:2012jc}. Other exotic explanations include \\cite{Yang:2012qi,Lawson:2012zu,Cline:2012hb}.  In this paper, we revisit the determination of the isotropic radio background from radio maps. An attempt in the same direction has been recently undertaken by Ref.~\\cite{Subrahmanyan:2013eqa}. Here we consider a more realistic and sophisticated Galactic description, we increase the number of maps employed, and adopt a different statistical technique.  In Sect. \\ref{sec:Surveys}, we present the maps which are more relevant for the isotropic background study. Sect. \\ref{sec:Models} describes how the Galactic emission is accounted for. The emission observed in radio skymaps from few MHz to hundreds of GHz is the result of diverse emission mechanisms and sources. The Milky Way activity accelerates CR electrons and positrons and amplifies magnetic fields, and a large fraction of the Galactic radio emission is produced via synchrotron radiation. In addition, thermal bremsstrahlung coming from warm and hot gas in the interstellar medium can contribute at such frequencies but is more localized nearby the Galactic plane. Discrete sources also provide an important fraction of the total emission. Finally, the extragalactic background from unresolved sources can dominate at high-latitudes.  We adopt a 4D modeling (3 spatial dimensions plus momentum) of CR synchrotron radiation, solving the CR transport equation by means of the publicly available code GALPROP~\\cite{Strong:1998pw}\\footnote{We used the GALPROP version 54.1.984, downloaded from Ref. \\cite{galpropweb}. The implementation in GALPROP of the magnetic field model employed in this analysis, as well as the theoretical and observational maps shown in this work, are available at \\cite{mapsweb} or by writing to taoso@cea.fr or regis@to.infn.it }. The propagation models are tuned to fit local CR spectra and gamma-ray data~\\cite{Strong:2010pr}. We employ a detailed description of the magnetic field following Refs.~\\cite{Jansson:2012rt,Jansson:2012pc}. Discrete sources are either masked or described by means of a template as discussed in Sect. \\ref{sec:Fit}, where we also describe the fitting procedure. The resulting estimates of the isotropic radio background is presented in Sect. \\ref{sec:Results}. Caveats and consequences of the analysis are discussed in Sect. \\ref{sec:Discussion}, while Sect. \\ref{sec:Conclusions} concludes. ",
        "conclusions": "\\label{sec:Conclusions}  \\begin{table}[t] \\begin{center} \\begin{tabular}{|c|c|c|c|c|} \\hline Frequency [MHz]& $T_E$ [K] & $T_{NC}$ [K] & zero-level [K] & $T_{\\rm cold\\,spot}$ [K]\\tabularnewline \\hline 22 & $(1.04 \\pm 0.24) \\times 10^4$ & $6.92 \\times 10^3$  & 5000 & $1.80\\times 10^4$ \\tabularnewline \\hline 45 & $(2.95 \\pm 0.34) \\times 10^3$ & $1.0 \\times 10^3$  & 544 & $3.84\\times 10^3$ \\tabularnewline \\hline 408 & 11.8 $\\pm 1.1$  & 2.61  & 3 & 12.14 \\tabularnewline \\hline 820 & 2.21 $\\pm 0.39 $ & 0.39  & 0.6 & 2.91 \\tabularnewline \\hline 1420 & 0.580 $\\pm 0.025$  & 0.09  & 0.2 & 0.589 \\tabularnewline \\hline 2326 & 0.073 $\\pm 0.013 $  & 0.024  & 0.08 & 0.098 \\tabularnewline \\hline \\end{tabular} \\end{center} \\caption{Estimates of the isotropic component $T_E$ at different frequencies from the best-fit model in our analysis, compared to the estimates $T_{NC}$ obtained from number counts, the zero-level error of the surveys, and the temperature $T_{\\rm cold\\,spot}$ of the coldest spot ($15^\\circ$ of radius) in the maps.} \\label{tab:results} \\end{table}   We collected the most complete observational sky maps at radio frequencies ranging from 22 MHz to 2326 MHz. The observed temperature is theoretically modeled by means of four components: Galactic diffuse synchrotron radiation, Galactic thermal bremsstrahlung, single sources, and isotropic emission. We employed different methods to assess the main uncertainties associated to Galactic modeling and to the treatment of source through masks or templates. The synchrotron emission is the main Galactic component and has been computed adopting a purely diffusive model for CRs (solving the transport equation with the GALPROP code~\\cite{Strong:1998pw}) and a magnetic field based on Refs. \\cite{Jansson:2012rt,Jansson:2012pc}. We considered different cases for the parameters defining its latitude profile (such as the vertical scale of the diffusion box $L$ and the $z$-behaviour of the magnetic field). Our main focus has been the determination of the isotropic background temperature $T_E$, and we found fully compatible results among the different methods, with a moderate scatter, which makes us to believe that the estimate is robust.  Our findings on $T_E$ lie above expectations from what can be estimated by extrapolation to low intensities of number counts of observed extragalactic sources: the reconstructed values of $T_E$ exceed them by a factor $\\gtrsim 3$ in the maps with the largest fraction of sky available (45 MHz, 408 MHz, and 1.4 GHz), which is where our estimates are more solid. Therefore, our results confirm the so called ``ARCADE'' excess, although our estimates are somewhat smaller than the results in Ref. \\cite{Fixsen:2009xn}.  In the radio maps, the isotropic component dominates the emission at high latitudes making the total profile quite flat at $b\\gtrsim 50^\\circ$. A Galactic component with a spatial distribution related either to the disc, or to a DM halo (as in the case of annihilating or decaying DM), or being very local would introduce some latitude dependence which is not seen in the data. For this reason, a Galactic origin for the mismatch between our estimate and the integral of number counts appears unlikely, unless some new model, with quite different properties with respect to what is usually considered, will be developed.  On the other hand, from statistical estimates of the source count at fluxes below the faintest sources that can be counted individually, no significant deviation from the standard extrapolation of extragalactic source count is expected above few hundreds of nJy~\\cite{Vernstrom:2013vva}.  We therefore conclude that observational radio maps point towards a puzzling excess. It could be due to to a systematic offset in the surveys (although in the same positive direction for all the maps), to the need of a profound modification of our current understanding of the Galactic synchrotron emission, or to a novel population of very faint extragalactic sources. The extragalactic interpretation is quite interesting, since its origin could be linked to a dark matter emission \\cite{paper1}.  Next future telescopes, and in particular LOFAR~\\cite{LOFAR} for the Galactic diffuse emission and SKA~\\cite{SKA} (with its precursors) for the source counts, will bring important tiles to solve this intriguing puzzle.   {"
    },
    "1402/1402.2504_arXiv.txt": {
        "abstract": "Sub-parsec binary massive black holes (BBHs) are long anticipated to exist in many QSOs but remain observationally elusive. In this paper, we propose a novel method to probe sub-parsec BBHs through microlensing of lensed QSOs. If a QSO hosts a sub-parsec BBH in its center, it is expected that the BBH is surrounded by a circum-binary disk, each component of the BBH is surrounded by a small accretion disk, and a gap is opened by the secondary component in between the circum-binary disk and the two small disks. Assuming such a BBH structure, we generate mock microlensing light curves for some QSO systems that host BBHs with typical physical parameters. We show that microlensing light curves of a BBH QSO system at the infrared-optical-UV bands can be significantly different from those of corresponding QSO system with a single massive black hole (MBH), mainly because of the existence of the gap and the rotation of the BBH (and its associated small disks) around the center of mass. We estimate the half-light radii of the emission region at different wavelengths from mock light curves and find that the obtained half-light radius vs. wavelength relations of BBH QSO systems can be much flatter than those of single MBH QSO systems at a wavelength range determined by the BBH parameters, such as the total mass, mass ratio, separation, accretion rates, etc.  The difference is primarily due to the existence of the gap.  Such unique features on the light curves and half-light radius-wavelength relations of BBH QSO systems can be used to select and probe sub-parsec BBHs in a large number of lensed QSOs to be discovered by current and future surveys, including the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS), the Large Synoptic Survey telescope (LSST) and Euclid. ",
        "introduction": "\\label{sec:intro}  Massive binary black holes (BBHs) on sub-parsec scales are unavoidable products of mergers of galaxies if each galaxy contains a central massive black hole \\citep[MBH; e.g.,][]{BBR80,Yu02, MM05}. In the $\\Lambda$CDM cosmology, big galaxies are formed through hierarchical mergers of small galaxies, therefore BBHs may exist in the centers of many galaxies. In gas-rich mergers, the consequent sub-parsec BBHs may be surrounded by a circum-binary disk, within which a gap is opened by the inspiraling secondary MBH, and each component of the BBH may be associated with a small accretion disk within the gap, fed by material infalling from the inner edge of the circum-binary disk \\citep[e.g.,][]{AL94, Ivanov99, Escala05, Hayasaki07, Hayasaki08, Dotti07, Cuadra09, DOrazio12, Rafikov12, Farris13}. These active BBHs may also emerge as QSOs but may have some unique features distinguishable from single active MBH systems.  Observational search for and identifying BBHs is important because they are not only probes to the current galaxy formation model but also main sources for gravitational waves. Recent observations have already revealed a few active MBH pairs (dubbed as dual AGNs) on kpc scale through various techniques \\citep[e.g.,][]{Komossa03, Wang09, Comerford09, Liu10, Shen11, Fu11, Rosario11, Yu11}, which are presumably the precursors of sub-parsec BBHs. However, spatially resolving and identifying sub-parsec BBHs (both active and inactive ones) is still beyond the capabilities of the present-day telescopes \\citep[e.g.,][]{Yu02}.  A number of methods have thus been proposed and applied to probe the BBH systems indirectly. Such methods include the double-peaked or asymmetric broad emission line \\citep[e.g.,][]{BL09, Tsalmantza11, Eracleous12, Ju13, Shen13}, flux ratio of different emission lines \\citep{Montuori11, Montuori12}, periodic variation of the QSO light curves \\citep[e.g.,][]{Valtonen08, Sesana12, GM12, Hayasaki12}, double broad relativistic Fe K$\\alpha$ lines \\citep[e.g.,][]{YL01, Sesana12, McKernan13}, and kinematic signature \\citep{Meiron13}, etc. Most of these methods depend on the orbital motions of BBHs. A number of sub-parsec BBH candidates have been reported but few of them were confirmed because of various complications in these methods, which make the differences of the properties of BBH systems from those of its alternatives ambiguous.  It is therefore imperative to devise new method/approach to unambiguously reveal sub-parsec BBHs/binary QSOs.  It has been demonstrated that the microlensing magnification of lensed QSOs \\citep[e.g.,][]{WPS90,WMS95} is a powerful tool to probe the size and even temperature structure of accretion disks, as the Einstein radius of a single star in the foreground lens galaxies is comparable to the disk size of background QSOs and the QSO flux variation is determined by the ratio of the Einstein radius to the size of the emission region \\citep[e.g.,][]{Wambsganss06}.  This method has been applied to estimate the disk sizes for more than a dozen lensed QSOs, and the results show that the estimated disk temperature profile is consistent with the thin disk model though the estimated disk size may be a factor of $\\sim 4-5$ larger than the standard thin disk prediction \\citep[e.g.,][]{WP91, Mortonson05, Eigenbrod08, Poindexter, Morgan10, Blackburne11,JV12}.  \\citet{DA11} and \\citet{AS12} suggest that alternative disk structure models, such as inhomogeneous disk or super-Eddington accretion disk, may be required to explain the microlensing results. Long term monitoring of microlensing of lensed QSOs are helpful to identify the detailed structure of accretion disks associated with these lensed QSOs.  The structure of disk accretion onto BBHs is fundamentally different from that onto a single MBH. The most distinct feature of an active BBH system should be the gap opened by the inspiraling secondary MBH \\citep[e.g.,][]{AL94, Ivanov99, Dotti07, Hayasaki07, Hayasaki08, Cuadra09}.  Direct detection of such a feature is not possible currently. In this paper, we propose a novel method to probe sub-parsec BBHs/binary QSOs by detecting signatures of the gap through the microlensing of lensed QSOs.  This paper is organized as follows. In section~\\ref{sec:config}, we first illustrate the configuration of an active BBH system and calculate the surface brightness distribution of such a system at the given UV-optical-infrared bands by adopting a simple thin disk accretion model. In section~\\ref{sec:mocklc}, mock microlensing light curves are generated by convolving the brightness distribution of background sources, either an active MBH or an active BBH system, with the foreground magnification maps. Adopting the Monte-Carlo microlensing analysis method developed by \\citet{Kochanek04}, we demonstrate that the size of emission region can be estimated for a BBH system through the mock light curves by assuming a single standard accretion disk model in section~\\ref{sec:disksize}. As shown in section~\\ref{subsec:results}, the expected relationship between disk size and wavelength for a BBH system show some apparent anomaly at some wavelength compared with that for a single MBH system. We propose that this anomaly can be taken as a signature of BBHs and should be useful for selecting BBH candidates from lensed QSOs with current and future facilities.  Discussions and conclusions  are given in sections~\\ref{sec:complications} and \\ref{sec:discon}. ",
        "conclusions": "\\label{sec:discon}  In this paper, we propose a novel method to probe sub-parsec BBH QSOs through the microlensing of lensed QSOs. If a QSO hosts a sub-parsec BBH in its center, it is expected that the BBH is surrounded by a circum-binary disk at outside region, each component of the BBH is surrounded by a small accretion disk, and a gap opened by the small component in between the circum-binary disk and the two small disks. Assuming such a structure for some hypothetical BBH systems, we generate mock microlensing light curves for BBH QSO systems with typical physical parameters. We show that the microlensing light curves of a BBH QSO system at some given bands can be significantly different from those of a single MBH QSO system because of the existence of the gap and the rotation of the BBH and its associated small disks around the center of mass. We estimate the half-light radius-wavelength relations from those mock light curves and find that the obtained half-light radius-wavelength relations of BBH QSO systems can be much flatter than that of single MBH QSO systems at a wavelength range determined by the BBH parameters, such as the total mass, mass ratio, separation, etc., which is primarily due to the existence of the gap. Such a unique feature of BBH QSO systems can be used to select and probe sub-parsec BBHs in a large number of lensed QSOs to be discovered in the near future.  Note that hints for the existence of BBHs in some QSO systems may be revealed by their SEDs \\citep[][see section~\\ref{subsec:diskmodel}]{GM12} and other spectral features (see the introduction), but these features may be model dependent and somewhat obscure due to complications in the emission processes. The new method proposed in this paper can provide stronger physical constraints on BBH systems and is less model dependent, as the structure of the disk accretion associated with BBH systems can be well resolved by microlensing. The signatures are particularly striking when the rotation effects are present (see Figures~\\ref{fig:f7} and \\ref{fig:f8}).  Many of the lensed QSOs may be single MBH systems, but some of them might be accretion systems with BBHs in their centers. The occurrence rate of BBHs in QSOs is an important quantity to determine the efficiency of using microlensing to probe sub-parsec BBHs in lensed QSOs. In principle, the BBH occurrence rate is determined by: (1) the active level of the two MBHs within the gap; and (2) the residence time of BBHs at the separation that has the most significant signals (i.e., $\\sim 500-1000\\rg$).  Most recent numerical simulations suggest that significant accretion can be induced onto the two MBH components although a low-density cavity (gap) at the center of the circum-binary disk is maintained \\citep[e.g.,][]{Hayasaki07, MM08, Roedig11, Noble12, Shi12, DOrazio12}. In this paper, we assume that the level of activities of sub-parsec BBH systems embedded in circum-binary disks are more or less the same as that of normal bright QSOs, therefore, the occurrence rate of BBH systems in QSOs is mainly determined by the BBH residence time. For those BBHs with total mass in the range of $10^7 -10^9\\msun$, mass ratio $q\\sim 0.1-1$ and separation $\\sim 500-1000 \\rg$, the residence time is roughly in the range from a few $10^5$ to a few $10^6$ yr, according to \\citet[][see Figures 1 and 3 therein]{Haiman09}.  The lifetime of normal QSOs is estimated to be in the range of a few times $10^7$ to $10^8$~yr from various methods \\citep[e.g.,][]{HH01, MW01, Jakobsen03, YL04, YL08, Marconi04, WW06, Worseck07, GSP08, Shankar09, Shankar13, LY11}.  Thus the occurrence rate of BBHs in lensed QSOs should be in the range of a few thousandth to a few percent.  The total number of currently known lensed QSOs is $\\sim 100$ (\\citealt{Kochanek06}; http://www.cfa.harvard.edu/castles; http://www-utap.phys.s.u-tokyo.ac.jp/\\~sdss/sqls/).  According to the BBH occurrence rate estimated above, at most, several of these lensed QSOs could host BBH systems with separation $\\la 1000\\rg$, which may be detectable through the microlensing event(s).  Ground-based surveys, such as the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS; which is already operating), the planned Large Synoptic Survey telescope (LSST) and Euclid, will monitor a large area of sky in multi-bands from UV, optical to infrared for multiple times over a long period (e.g., $10$~yr). It is expected that the Pan-STARRS and the LSST will discover $\\sim 3000-8000$ luminous QSOs that are gravitational lensed into multiple images by foreground galaxies \\citep{Abell09, OM10}. It is possible that tens to one hundred of these lensed QSOs contain BBHs in their centers according to the BBH occurrence rate estimated above.  Long term monitoring of these lensed QSOs by the Pan-STARRS, LSST, Euclid and other future facilities may provide a powerful and efficient way to discover sub-parsec BBHs unambiguously through the anomaly in the $r_{1/2}-\\lambda$ relation estimated from the QSO light curves."
    },
    "1402/1402.0392_arXiv.txt": {
        "abstract": "{Simple molecules such as H$_2$CO and CH$_3$OH in protoplanetary disks are the starting point for the production of more complex organic molecules. So far, the observed chemical complexity in disks has been limited because of freeze-out of molecules onto grains in the bulk of the cold outer disk.} {Complex molecules can be studied more directly in transitional disks with large inner holes because these have a higher potential of detection through the UV heating of the outer disk and the directly exposed midplane at the wall.}{We used Atacama Large Millimeter/submillimeter Array (ALMA) Band 9 ($\\sim$680 GHz) line data of the transitional disk Oph IRS 48, which was previously shown to have a large dust trap, to search for complex molecules in regions where planetesimals are forming.}{We report the detection of the H$_2$CO 9(1,8)--8(1,7) line at 674 GHz, which is spatially resolved as a semi-ring at $\\sim60$ AU radius centered south from the star. The inferred H$_2$CO abundance is $\\sim 10^{-8}$ , derived by combining a physical disk model of the source with a non-LTE excitation calculation. Upper limits for CH$_3$OH lines in the same disk give an abundance ratio H$_2$CO/CH$_3$OH $>$0.3, which indicates that both ice formation and gas-phase routes play a role in the H$_2$CO production. Upper limits on the abundances of H$^{13}$CO$^+$, CN and several other molecules in the disk were also derived and found to be consistent with full chemical models.}{The detection of the H$_2$CO line demonstrates the start of complex organic molecules in a planet-forming disk. Future ALMA observations are expected to reduce the abundance detection limits of other molecules by 1--2 orders of magnitude and test chemical models of organic molecules in (transitional) disks.} ",
        "introduction": "Planets are formed in disks of dust and gas that surround young stars. Although the chemical nature of the gas is simple, with only small molecules such as H$_2$, CO, HCO$^+$ , or H$_2$CO detected so far, the study of molecules in protoplanetary disks has resulted in a much better understanding of the origin of planetary systems \\citep[e.g.][]{WilliamsCieza2011,Henning2013}. Molecular line emission serves as a probe of disk properties, such as density, temperature, and ionization. Furthermore, simple species are the start of the growth of more complex organic and possibly prebiotic molecules \\citep{Ehrenfreund2000,Mumma2011}.  Molecules in disks are incorporated into icy planetesimals that eventually grow to comets and asteroids that may have delivered water and organic material to Earth. Therefore, a better understanding of the chemical composition of protoplanetary disks where these icy bodies are formed provides insight into the building blocks of comets and Earth-like planets elsewhere in the Universe.  Protoplanetary disks have sizes of up to a few 100 AU, which makes them similar to or larger than our own solar system ($\\sim$50 AU radius). However, at the distance of the nearest star-forming regions these disks subtend less than a few arcsec on the sky, so that telescopes with high angular resolution and high sensitivity are needed to study their chemical composition. Most of the disks observed so far show no high chemical complexity \\citep[e.g.][]{Dutrey1997,Thi2004,Kastner2008,Oberg2010}. In the surface layers of the disk, molecules are destroyed by photodissociation by ultraviolet (UV) radiation from the central protostar. In the outer disk and close to the disk midplane, temperatures quickly drop to 100 K and lower, where all detectable molecules, including CO, freeze out onto dust grains at temperatures determined by their binding energies \\citep{Bergin2007}. The chemical composition thus remains hidden in ices. Only a very small part of the ice molecules is brought back into the gas phase by nonthermal processes such as photodesorption.  Transition disks have a hole in their dust distribution and thus form a special class of protoplanetary disks \\citep{WilliamsCieza2011}. The hole allows a view into the usually hidden midplane composition because the ices at the edge of this hole are directly UV irradiated by the star, which results in increased photodesorption and thermal heating of the ices \\citep{Cleeves2011}. The hole is also an indicator that the disk may be at the stage of forming planets. So far, the chemical composition of the outer regions of transitional disks appears to be similar to that of full disks, with detections of simple molecules including H$_2$CO \\citep{Thi2004,Oberg2010,Oberg2011}. However, these data were taken with a typical resolution of $>$2''.   The study of the chemistry in protoplanetary disks has recently gained much more perspective due to the impending completion of the Atacama Large Millimeter/submillimeter Array (ALMA). ALMA allows us to study astrochemistry within protoplanetary disks at an unprecedented level of complexity and at very small scales. It has the sensitivity to detect not only the dust, but also the gas inside dust gaps in transitional disks  \\citep{vanderMarel2013,Casassus2013,Fukagawa2013,BruderervdM}.  The disk around Oph IRS 48 forms a unique laboratory for testing basic chemical processes in planet-forming zones. IRS 48 is a massive young star ($M_* = 2 M_{\\odot}$, $T_* \\sim$10 000 K) in the Ophiuchus molecular cloud (distance 125 pc) with a transition disk with a large inner dust hole, as revealed by mid-infrared imaging, which traces the hot small dust grains \\citep{Geers2007}. The submillimeter continuum data (685 GHz or 0.43 mm) from thermal emission from cold dust obtained with ALMA show that the millimeter-sized dust is concentrated on one side of the disk, in contrast to the gas and small dust grains. Gas was detected inside the dust gap down to 20 AU radius with ALMA data \\citep{BruderervdM}, and strong PAH emission is also observed from within the cavity \\citep{Geers2007}. The continuum asymmetry has been modeled as a major dust trap \\citep{vanderMarel2013}, triggered by the presence of a substellar companion at $\\sim$20 AU. The dust trap provides a region where dust grains concentrate and grow rapidly to pebbles and then planetesimal sizes, producing eventually what may be the analog of our Kuiper Belt.  \\citet{BruderervdM} presented and modeled the ALMA CO and continuum data, together with complementary data at other wavelengths, to derive a three-dimensional axisymmetric physical model of the IRS 48 disk. One important conclusion is that the dust in the disk is warm, even out to large radii, because the UV radiation can pass nearly unhindered through the central hole. This implies that the dust temperature is higher than 20 K throughout the disk so that CO does not freeze out close to the midplane of the disk \\citep{Collings2004,Bisschop2006}. The lack of a freeze-out zone of CO is important, since much of the chemical complexity in a protoplanetary disks is thought to start with the hydrogenation of the CO ice \\citep{Herbst2009}. In the absence of CO ice, gas-phase chemistry is the current main contributor to complex molecule formation. Alternatively, the disk may have been colder in the past before the hole was created, with the ices produced at that time now evaporating. One of the simplest complex organic molecules, H$_2$CO, can form both through hydrogenation of CO-ice \\citep{Tielens1982,Hidaka2004,Fuchs2009,Cuppen2009} and through gas-phase reactions.  H$_2$CO has been detected in several astrophysical environments, such as the warm inner envelopes of low- and high-mass protostars and protoplanetary disks including transitional disks \\citep{Dutrey1997,Ceccarelli2000,Aikawa2003,Thi2004,Bisschop2007,Oberg2010,Oberg2011,Qi2013}. In contrast with H$_2$CO, CH$_3$OH can only be formed through ice chemistry \\citep{Geppert2006,Garrod2006FD}, which means that the H$_2$CO/CH$_3$OH ratio gives information on the H$_2$CO formation mechanism. Furthermore, H$_2$CO is a very interesting molecule for comparing the chemical composition of disks, comets, and our solar system, because H$_2$CO- and CN-bearing molecules such as HCN and CN are precursors of amino acids. Synthesis of amino acids occurs in large asteroids in the presence of liquid water, for instance, through the Strecker synthesis route \\citep{Ehrenfreund2000}.   In this work we present the detection of the H$_2$CO 9(1,8)--8(1,7) line in IRS~48 down to scales of $\\sim$30 AU, a high-excitation line originating from a level at 174 K above ground. We model its abundance and discuss the implications of the origin of this molecule in combination with the nondetection of other molecular lines. We were only able to detect this line thanks to the tremendous increase in sensitivity at these high frequencies ($\\sim$ 670 GHz) using ALMA. ",
        "conclusions": "We observed the Oph IRS 48 protoplanetary disk with ALMA Early Science at the highest frequencies, around 690 GHz, allowing the detection of warm H$_2$CO and upper limits on the abundances of several other molecules including CH$_3$OH, H$^{13}$CO$^+$ , and CN lines at unprecedented angular resolution. \\begin{enumerate}  \\item We detected and spatially resolved the warm H$_2$CO 9(1,8)-8(1,7) line, which reveals a semi-ring of emission at $\\sim$60 AU radius centered south from the star. No emission is detected in the north. This demonstrates that H$_2$CO, an ingredient for building more complex organic molecules, is present in a location of the disk where planetesimals and comets are currently being formed.  \\item The H$_2$CO emission was modeled using a physical disk model based on the dust continuum and CO emission \\citep{BruderervdM}, using three different trial abundance profiles. None of the profiles were able to match the observed data exactly, but the absolute flux indicates an abundance with respect to H$_2$ of $\\sim10^{-8}$.  \\item The combination of the H$_2$CO abundance in combination with upper limits for the CH$_3$OH emission indicates a H$_2$CO/CH$_3$OH ratio $>$0.3. This limit together with the overall abundance suggests that both solid-state and gas-phase processes occur in the disk.  \\item Although the H$_2$CO emission is located only on the southern side of the disk, just like the millimeter dust continuum, the offset with the continuum peak and the low $S/N$ do not allow a firm claim on a relation with the dust-trapping mechanism.  \\item The upper limit for H$^{13}$CO$^+$ indicates an HCO$^+$ abundance of $<10^{-8}$, consistent with our model. The upper limit for CN of $10^{-7.3}$ relative to H$_2$ is directly at the level of that predicted by our model. Upper limits on the abundances of the other targeted molecules are consistent with earlier observations.  \\item Future ALMA observations of intrinsically stronger lines will allow abundances to be measured that are one or more orders of magnitude below the upper limits derived here. This will allow full tests of the chemistry of simple and more complex molecules in transitional disks. \\end{enumerate}"
    },
    "1402/1402.7098_arXiv.txt": {
        "abstract": "Observations indicate that most massive galaxies contain a supermassive black hole, and theoretical studies suggest that when such galaxies have a major merger, the central black holes will form a binary and eventually coalesce. Here we discuss two spectral signatures of such binaries that may help distinguish them from ordinary AGN.   These signatures are expected when the mass ratio between the holes is not extreme and the system is fed by a circumbinary disk.   One such signature is a notch in the thermal continuum that has been predicted by other authors; we point out that it should be accompanied by a spectral revival at shorter wavelengths and also discuss its dependence on binary properties such as mass, mass ratio, and separation. In particular, we note that the wavelength $\\lambda_n$ at which the notch occurs depends on these three parameters in such a way as to make the number of systems displaying these notches  $\\propto \\lambda_n^{16/3}$; longer wavelength searches are therefore strongly favored. A second signature, first discussed here, is hard X-ray emission with a Wien-like spectrum at a characteristic temperature $\\sim 100$~keV produced by Compton cooling of the shock generated when streams from the circumbinary disk hit the accretion disks around the individual black holes.  We investigate the observability of both signatures.   The hard X-ray signal may be particularly valuable as it can provide an indicator of black hole merger a few decades in advance of the event. ",
        "introduction": "Binary supermassive black holes are objects of great intrinsic interest. That two massive black holes should coexist in a single galaxy appears to be a natural corollary of the standard theory of galaxy formation, in which galaxies are built up by mergers from smaller galaxies (see \\citealt{2012RAA....12..917S} for a recent review), and the observation that nearly every galaxy whose luminosity is more than a fraction of the characteristic galaxy luminosity $L_*$ contains a supermassive black hole in its nucleus \\citep{KormendyHo2013}. Their progress toward ultimate merger, however, depends upon the gas content and stellar orbital distribution of the galaxy in which they live \\citep{BBR80,GouldRix2000,2003ApJ...596..860M,2004ApJ...607..765E,2005ApJ...630..152E,2006ApJ...642L..21B,Cuadra2009,2009JPhCS.154a2049P,2009ApJ...695..455B,2011ApJ...732...89K,Vasiliev13,2013ApJ...773..100K}. In addition, binary supermassive black holes should be prodigious sources of gravitational wave emission during their actual merger, and detectable by pulsar timing arrays even well before merger (see \\citealt{2013MNRAS.433L...1S} for a recent calculation).  Despite considerable effort, observational examples are few and limited in their character.  There are now a number of ``duals\", which are two supermassive black holes that are in the same galaxy but are not gravitationally bound to each other in a true binary \\citep{Komossa03,Comerford13,Liu13,Woo13}.   These can be identified through a number of techniques, including X-ray or optical imaging revealing two AGN-like point-sources in a single galaxy and doubled line profiles.   These methods require that both black holes receive a large enough accretion flow to ``light up\" as AGN.  Analogous spectral methods have been employed to search for genuine binaries, but they are difficult to use.   For example, if both members of the black hole binary are AGN, and their separation is large enough that their mutual orbital speeds are small compared to the width of broad emission lines, the offset between their respective broad line profiles will be small compared to their intrinsic widths and hard to discern; conversely, if the separation is small enough to produce a larger orbital velocity, the broad emission line regions of the two will overlap, and the lines will respond to the combined gravitational potential and ionizing radiation of both black holes \\citep{ShenLoeb10,2012ApJS..201...23E}.    Alternatively, shifts in broad emission line peaks or centroids may perhaps indicate binaries \\citep{Loeb13,Decarli13}.  \\cite{ArtymLubow94} showed that circumbinary disks around binary systems with mass ratios $q \\gtrsim 4 (h/r) \\alpha^{1/2}$ have low density cavities around the binary because stable circular orbits do not exist within a radius $\\sim 2a$ of their centers of mass (here $q = M_2/M_1$ for $M_1$ the larger mass and $M_2$ the smaller; $h/r$ is the disk aspect ratio; $\\alpha$ is the ratio of stress to pressure within the disk; and $a$ is the binary's semimajor axis).    In the context of supermassive black hole binaries, the lower bound on $q$ for creating a cavity can be quite small. Although initial studies of circumbinary disk dynamics suggested that little of the mass accreting through the outer regions of the disk would pass through the inner edge at $r \\simeq 2a$ and be captured by the binary \\citep{Pringle91}, recent work has indicated that the ``leakage fraction\" is at least $\\sim 10\\%$ \\citep{MM08,Shi2012,Roedig12,Farris13} and possibly unity (Shi et al. 2014 in preparation).  All studies so far that have examined the question find that the majority of the accretion goes to the secondary rather than the primary when the mass ratio is not too small \\citep{Roedig12,Farris13}. This implies that accretion will increase the mass ratio. It also implies that if the accretion rate fed to the circumbinary disk is enough to support a luminous AGN, as is likely in a galaxy that is the product of a recent merger \\citep{Hopkins12,Mayer13}, at least one and possibly both of the black holes will acquire its own accretion disk and appear as an AGN.  There has been much interest in finding distinctive electromagnetic signatures of these systems, partially to supplement future gravitational wave observations, but also so that they may be found and studied even before instruments capable of detecting their gravitational waves are built.   Several papers have noted that the absence of optically thick gas in the low-density region between the circumbinary disk and the accretion disks around each black hole (from here on, called ``minidisks\") will produce a dip in the thermal continuum spectrum \\citep{Roedig12,Tanaka12,GM12,Kocsis12,Tanaka13}.    However, all but  one \\citep{Roedig12} suppose that minidisks either do not exist or receive matter at a rate substantially below the accretion rate in the circumbinary disk, thus underestimating the thermal continuum at frequencies above the dip, nor has there been any study of the summed spectrum's dependence on parameters. More importantly, all previous work (except \\cite{Tanaka13b}, discussed later) neglects the spectral contribution from the energy released when the fluid streams that carry matter on ballistic orbits from the inner edge of the circumbinary disk to the outer edges of the individual disks shock and release their kinetic energy as their contents join those disks.  Here we present a more complete account of these radiation processes and explore how they depend on the binary parameters $M=M_1 + M_2$, $q$, and $a$.  We also use estimates of binary lifetimes at their different evolutionary stages to predict which signals should be most common and easiest to observe.  In Section~2 we discuss thermal radiation from the circumbinary disk and the two minidisks.  In Section~3 we discuss the previously neglected radiation from the hot spots produced by impact of the ballistic streams onto the minidisks.  In Section~4 we explore the observability of our signatures as a function of masses, mass ratios, and separations, and we present our conclusions in Section~5. ",
        "conclusions": "Previous work pointed out that the interruption of a circumbinary disk surrounding a pair of black holes leads to a cut-off in its thermal continuum spectrum at much lower energies than the cut-off expected in a disk around a single black hole of the same mass.   We have shown that the spectrum should, in general, revive at energies $\\sim 10\\times$ greater because the minidisks surrounding the individual black holes become bright at those energies.   In addition, we have, for the first time, discussed how the properties of such a spectral ``notch\" depend upon system mass, mass ratio, and separation.  We have also introduced a new continuum spectral signature characteristic of black hole binaries that will be most prominent at rest-frame energies of tens to hundreds of keV.  This radiation comes from the hot spots produced when accretion streams fall nearly ballistically from the inner edge of the circumbinary disk and hit the outer edges of the minidisks around the individual supermassive black holes.  These hot spots should radiate Wien spectra with temperatures $\\sim 100$~keV, although given the current limitations of hard X-ray detectors, these spectra might be easier to detect as spectral hardening at tens of keV or above, or possibly by a periodic modulation in the X-ray flux at those energies.  There are, however, certain difficulties associated with searches for either one. The number of sources in the sky with a notch in the thermal continuum spectrum at wavelength $\\lambda_*$ is extremely sensitive to $\\lambda_*$ itself ($\\propto \\lambda_*^{16/3}$) as well as the redshift of the source ($\\propto (1+z)^{-13/3}$) and even the portion of the notch falling at $\\lambda_*$ (there are $\\sim 200\\times$ more systems at any given time for which the minimum of the notch is at $\\lambda_*$ than those for which the low-frequency edge of the notch is at $\\lambda_*$). If these searches wish to use the high sensitivity of optical band instruments, they will be most sensitive to systems in which the parameter combination $\\dot m M_8^{-1} (a/100R_g)^{-3} (1+z)^{-4} \\sim 10^{-3}$. Similarly, for the Wien bump of the hot-spots to stand out clearly against the coronal hard X-rays generated by the usual processes in the innermost portions of the minidisks, it is necessary for $(a/100R_g) \\lesssim 1$, although it is possible that modulation of the hot-spot luminosity at a period comparable to the orbital frequency may relax this criterion somewhat.   Thus, these signatures are particularly powerful for identifying supermassive black hole binaries with separations $\\sim 10^{2\\pm 1}R_g$ or less.  In any event, both spectral signatures---the notch at longer wavelengths and the Wien bump in the hard X-ray band---are distinctive signatures of accretion onto black hole binaries.   It is hard to imagine another system that could produce either, much less both, of these features.    Thus, detection of either one would represent a strong suggestion, and detection of both a likely discovery of a supermassive binary black hole system on its way toward merger."
    },
    "1402/1402.3863_arXiv.txt": {
        "abstract": "We explore a cyclic universe due to phantom and quintessence fields. We find that, in every cycle of the evolution of the universe, the phantom dominates the cosmic early history and quintessence dominates the cosmic far future. In this model of universe, there are infinite cycles of expansion and contraction. Different from the inflationary universe, the corresponding cosmic space-time is geodesically complete and quantum stable. But similar to the Cyclic Model, the flatness problem, the horizon problem and the large scale structure of the universe can be explained in this cyclic universe. ",
        "introduction": "The theory of inflationary cosmology has become the standard paradigm for the early universe, which was firmly established in the past several decades.  Inflationary cosmology not only naturally solves the flatness problem, the horizon problem and the magnetic monopole problem, but also elegantly explains the large scale structure of the universe.  Despite of its success, it may be not perfect yet and suffer from some conceptual problems as follows.  First, a few common questions may be almost always asked as follows: What is the inflaton field? How to fine-tune the scalar potential in order that it is sufficiently flat?  How did the universe begin and why did it start to inflate? These questions are of great importance for the theory consummation. Second, inflationary cosmology suffers from the trans-Planckian problem \\cite{robert:12}. In inflationary cosmology, the fluctuations observed in the cosmic microwave background had wavelengths at the beginning of inflation smaller than the Planck scale. It is generally believed that the quantum gravity would become important at sub-Planckian scales. So the quantum gravity may produce different distribution on sub-Planckian wavelengths. This different distribution would be inflated and generate an uncertain correction to the predictions for the cosmic microwave background anisotropy. Third, Borde el al \\cite{borde:01} showed that the inflationary space-time is geometrically incomplete in the past directions of the history of the universe as the expansion rate averaged along the geodesic is positive, i.e., $\\left <H \\right> >0$. This incompleteness theorem relies on neither the energy condition nor the Einstein's equation. In order to solve these problems, a number of alternative scenarios, e.g., the cyclic universe \\cite{paul:02,baum:07}, the bounce universe \\cite{bounce:10} and the emergent universe \\cite{ell:046, ell:047, ell:048, ell:049, ell:0410, ell:0411} are proposed.  In the Cyclic Model \\cite{paul:02}, the universe undergoes an endless sequence of cosmic expansion and contraction. The cosmic temperature and energy density are all finite at each transition point from contraction to expansion or from expansion to contraction.  The Cyclic Model can solve not only the flatness problem, the horizon problem, the large scale structure problem, but also the Big-Bang singularity problem, the current cosmic accelerating problem and so on. Therefore, the cyclic model may present a complete description of all phases of the cosmic evolution.  The cyclic universe model \\cite{paul:02,baum:07} is originally constructed in the brane-world scenario. In this paper, we alternatively construct a cyclic universe by simply introducing a phantom matter and quintessence matter. In this cyclic universe, there are infinite cycles of expansion and contraction. Different from the inflationary universe, the corresponding cosmic space-time is geometrically complete and quantum stable. Similar to the Cyclic Model, the flatness problem, the horizon problem and the large scale structure of the universe can be explained in this cyclic universe.  This paper is organized as follows. In section~\\ref{sec:chameleon}, we present the energy density for the phantom and quintessence, respectively. Then, by taking into account both the relativistic  matter and dust matter,  we construct a cyclic model for the universe and investigate its cyclic evolution in section~\\ref{sec:model}. % The total cosmic energy density is zero at the transition points, either from expansion to contraction or from contraction to expansion. Therefore, the Big-bang and Big-crunch singularity problems are automatically vanishing. In section~\\ref{sec:geo}, we demonstrate that the cosmic space-time is geodesically complete for both time-like and null geodesics. The particle and the event horizon are all infinite in this space-time. So the horizon problem is solved. In section~\\ref{sec:entr}, we discuss the variation of entropy of the cyclic universe and the violation of second law of thermodynamics due to the presence of phantom matter. In section~\\ref{sec:quant}, using the Wheeler-DeWitt equation, we find that the {system} is quantum stable although a negative scalar potential is present. In section~\\ref{sec:scaleinvariant}, we verify that the primordial scale-invariant power spectrum problem could be generated in this cyclic model. Finally, conclusions and discussions are given in section~\\ref{sec:conclusion}.  Throughout this paper, we adopt the system of units in which $G=c=\\hbar=1$ and the metric signature $(-,\\ +,\\ +,\\ +)$. ",
        "conclusions": "\\label{sec:conclusion}  In conclusion, we have explored a cyclic model for the universe by simply considering a phantom field and a quintessence field. In every cycle of the cosmic evolution, the phantom field is dominant in the very early universe and the quintessence field is dominant in the far future.  There are infinite cycles of cosmic expansion and contraction in this model. At the turning points from expansion to contraction or from contraction to expansion, the cosmic energy density is zero. So the Hubble horizon is infinite at these {turning points}. We also find that the cosmic spacetime is geodesically complete for both time-like and null geodesics. Therefore, there is no Big-bang or Big-crunch singularity in the space-time. We calculate the particle horizon and the event horizon, respectively, and find they are infinite. So the horizon problem is vanishing. Using the Wheeler-DeWitt equation, we show the system is quantum stable although a negative scalar potential is present. Finally, we prove that the flatness of spacetime and the primordial scale-invariant power spectrum could be generated in this cyclic model which is consistent with the findings of Gratton et al \\cite{gra:04} and Boyel et al. \\cite{boy:06}."
    },
    "1402/1402.0321_arXiv.txt": {
        "abstract": "{We perform a combined analysis of the recent AMS-02 data on electrons, positrons, electrons plus positrons and positron fraction, in a self-consistent framework where we realize a theoretical modeling of all the astrophysical components that can contribute to the observed fluxes in the whole energy range. The primary electron contribution is modeled through the sum of an average flux from distant sources and the fluxes from the local supernova remnants in the Green catalog. The secondary electron and positron fluxes originate from interactions on the interstellar medium of primary cosmic rays, for which we derive a novel determination by using AMS-02 proton and helium data. Primary positrons and electrons from pulsar wind nebulae in the ATNF catalog are included and studied in terms of their most significant (while loosely known) properties and under different assumptions  (average contribution from the whole catalog, single dominant pulsar, a few dominant pulsars). We obtain a remarkable agreement between our various modeling and the AMS-02 data for all types of analysis, demonstrating that the whole AMS-02 leptonic data admit a self-consistent interpretation in terms of astrophysical contributions.}    \\begin{document} ",
        "introduction": "\\label{sec:intro} A huge experimental effort undertaken in the last decades has led to increasingly accurate measurements of cosmic rays (CRs) by means of space borne detectors. In particular, excellent data have been provided for the positron fraction ($e^+/(e^+ + e^-)$) and for the absolute positron ($e^+$),  electron ($e^-$) and total ($e^+ + e^-$) CR spectra by the Pamela \\citep{2009Natur.458..607A,2013arXiv1308.0133P,2011PhRvL.106t1101A} and {\\it Fermi}-LAT \\citep{2012PhRvL.108a1103A,2010PhRvD..82i2004A} Collaborations. Very recently, also the AMS-02 Collaboration has provided its first data on the positron fraction spectrum measured by the Alpha Magnetic Spectrometer installed on the International Space Station \\citep{Aguilar:2013qda}, and preliminary results for the other leptonic observables \\citep{electrons_AMS02,totalelectrons_AMS02}. The data cover an energy range spanning from about one GeV up to hundreds of GeV, depending on the species.  The most direct interpretation of these data refers to secondary production from nuclear collisions with primary CRs and the atoms of the interstellar medium (ISM), as well as to galactic astrophysical sources injecting fresh primary leptons into the ISM (see e.g. Ref. \\cite{2012APh....39....2S} and references therein, and Refs. \\cite{Hooper:2008kg,Profumo:2008ms,2009AA...501..821D,2009PhRvL.103h1104M,2010AA...524A..51D, 2013arXiv1312.3483D,2013APh....49...23E,Gaggero:2013nfa,Gaggero:2013oga,Gaggero:2013rya,Grasso:2013nyx,Yuan:2013eba,Cholis:2013lwa}). The observed raise of the positron fraction, firstly reported by the Pamela Collaboration and confirmed with higher precision by the AMS-02 data, has stimulated an extensive speculation on a possible dark matter (DM) contribution at high energies \\citep{2013arXiv1305.0084I,2013arXiv1307.6204B,2013arXiv1312.7841S,2013arXiv1309.2570I,2013PhRvD..88j3509D,2014PhLB..728...58H,Cirelli:2008pk,Zhao:2014nsa,Baek:2014awa,Kopp:2014tsa,Hryczuk:2014hpa,Choi:2013oaa,Ibarra:2013cra,Bergstrom:2013jra,Jin:2013nta,DeSimone:2013fia,Feng:2013zca,Masina:2013yea}.  In this paper, we explore at which level the new AMS-02 data on the whole leptonic observables may be accommodated in a purely astrophysical ($i.e.$ without invoking contributions from exotic sources such as DM annihilation) scenario, which counts the contribution from powerful stellar sources as well as from secondary reactions among primary CRs and atoms in the ISM. Specifically, we study the possible contribution that supernova remnants (SNRs) can give to high energy electron fluxes, and pulsar wind nebulae (PWN) to both positron and electron spectra. Whenever available, independent data on these sources are taken into account as constraints on their emission properties. We also employ the new preliminary AMS-02 data for proton and helium cosmic fluxes, in order to obtain a new evaluation of the secondary  $e^+$ and $e^-$ fluxes. All these components are added together and propagated in the Milky Way. A key point of our analysis is the requirement for any theoretical model to fit {\\it simultaneously} all the four AMS-02 leptonic observables, namely the $e^+/(e^+ + e^-)$, $e^+$, $e^-$ and $e^+ + e^-$ spectra. As we will show in next Sections, we find several astrophysical models compatible with all the AMS-02 leptonic data. Our results imply that a consistent and global picture of the Galaxy is possible, at least for the leptonic sector. Even more so, we show that high precision, low energy positron data, such as the ones collected by AMS-02, are on the verge of acting as a remarkable tool to constrain propagation models, and cooperate with the boron-to-carbon observations to this fundamental task. Our analyses are all implemented by a thorough estimation of the underlying possible uncertainties. This method allows us to derive  predictions on the AMS-02 data at very high energy, expected after some increase in the collected statistics. As a final remark, we notice that our analysis indicates that there is no particular need to invoke DM annihilation in the halo of the Milky Way in order to explain high energy positron and electron data. ",
        "conclusions": "\\label{sec:concl}  In this paper we have performed a combined analysis of the recent AMS-02 data on the electron flux, positron flux, electrons plus positrons flux and positron fraction, in a theoretical  framework that self-consistently accounts for all the astrophysical components able to contribute to the leptonic fluxes in their whole energy range.  The primary electron contribution is modeled through the sum of an average flux produced by distant sources and the fluxes arising from the local supernova remnants in the Green catalog. The secondary electron and positron fluxes originate from interactions on the interstellar medium of primary cosmic rays, for which we have derived a novel determination by using AMS-02 proton and helium data. Finally, the pulsar wind nebulae contribution to the positron (and electron) fluxes at high energies relies on a modeling of the sources reported in the ATNF catalog (where information on age, distance and total emitted power of the pulsars is available), under a number of  assumptions on the way the different pulsars might contribute to the local fluxes. We have in fact specifically performed three different type of analysis: we have studied the average contribution from the whole catalog; we have investigated if the ATNF catalog contains a single, dominant, pulsar which alone can allow agreement with the data; finally we have examined the possibility that a few powerful sources in the ATNF catalog may concurrently contribute to the local observed fluxes.  For all three different types of analysis, we obtain a remarkable agreement between our modelings and the whole set of AMS-02 data. The supernova remnants and the secondary contribution are able to properly explain the electron data and the low-energy part of the positron spectra, and to some extent they also point toward a disfavoring of small cosmic-rays confinement volumes. The high-energy part of the positron flux, which has received great attention because of its implications not only for the astrophysics of sources but also for dark matter studies, finds a remarkable solution in terms of pulsars present in the ATNF catalog. We find that AMS-02 data can be properly explained either in the case of an average contribution from the whole catalog, or for the situation where a single and close pulsar is the dominant contributor, or even in the case where a few and powerful dominant pulsars in the catalog are mostly contributing. For all cases, we have identified the required ranges of the relevant parameters (spectral index and efficiency of the emission) for the contributing pulsars, once the other parameters (age, distance, total emitted power) are fixed at their values reported in the ATNF catalog.    We can therefore conclude that the whole set of AMS-02 leptonic data admits a self-consistent interpretation in terms of astrophysical contributions.  Alternative solutions, like e.g. a dark matter production of electrons and positrons, are indeed a viable alternative or complementary possibility: however, a self-consistent solution in terms of purely astrophysical sources can be properly met."
    },
    "1402/1402.2748_arXiv.txt": {
        "abstract": "We investigate the possibility of arbitrarily high energy in the center of mass(CM) frame of colliding particles in the vicinity of the infinite red-shift surface of the spherically symmetric, static charged  regular black holes (Bardeen black hole, Ay\\'{o}n-Beato and Garc\\'{i}a black hole, and Hayward black hole). We show that the CM energy of colliding particles near the infinite red-shift surface of the \\emph{extreme} regular black holes are arbitrarily large while the non-extreme regular black holes have the finite energy. We also compute the equation of innermost stable circular orbit(ISCO), marginally bound circular orbit(MBCO) and circular photon orbit(CPO) of the above regular black holes, which are most relevant to black hole accretion disk theory. ",
        "introduction": "Recently, Ba\\~{n}ados, Silk and West(hereafter BSW) \\cite{bsw}  demonstrated that particles falling freely from rest exterior of a rotating extremal black hole can produce an infinite amount of high CM energy. In a semi-realistic setup, this energy could be higher than the Planckian energy, so that one might think extremal black hole could be act as super high energy particle accelerator.  Soon the appearance of this work in the literature several practical objections to the proposed mechanism has been raised. Particularly in \\cite{berti}, the authors have shown that there is an astrophysical bound i.e. maximal spin, back reaction effect and gravitational radiation etc. on that CM energy due to the Thorn's bound \\cite{thorn} i.e. $a=0.998M$ ($M$ is the mass and $a=\\frac{J}{M}$ is the angular momentum per unit mass of the black hole or Kerr parameter). Also Jacobson et al. \\cite{jacob} showed that CM energy in the near extremal situation for  Kerr black hole is $\\frac{E_{cm}}{2m_{0}}\\sim \\frac{2.41}{(1-a)^{1/4}}$. Lake \\cite{lake} found that the CM energy at the inner horizon of a static Reissner-Nordstr{\\o}m(RN) black hole and Kerr black holes are limited. Grib et al. \\cite{grib} investigated the CM energy using the multiple scattering mechanism. Also in \\cite{grib1}, the same authors computed the CM energy of particle collisions in the ergo-sphere of the Kerr black hole. The collision in the ISCO particles was investigated by the Harada et al. \\cite{harada} for Kerr black hole.  Liu et al. \\cite{liu} studied the BSW effect for Kerr-Taub-NUT(Newman, Unti, Tamburino) space-time and proved that the CM energy depends upon  both the Kerr parameter ($a$) and the NUT  parameter ($n$) of the space-time. Li et al. \\cite{li} proved that the Kerr-AdS  black hole space-time could be act as a particle accelerator. Studies were done by Zhong et al.\\cite{zong} for RN-de-Sitter  black hole and found that infinite energy in the CM frame near the cosmological horizon.  Said et al. \\cite{said} studied the particle accelerations and collisions in the back ground of a cylindrical black holes. In \\cite{piran}, the authors suggested that using collisional Penrose process the emitted massive  particles can only be gain $\\sim 30$ percentage of the initial rest energy of the in-falling particles.  In \\cite{wei1}, the authors discussed the CM energy for the Kerr-Newman black hole. For Kerr-Sen black hole, the CM energy is diverging  also discussed in \\cite{wei}. It was discussed in \\cite{zaslav} regarding the BSW process for spherically symmetric RN black hole.  Zhu et al. \\cite{zhu} showed that general stationary charged black holes as charge particle accelerators.  In \\cite{frolov}, the author demonstrated that a weakly magnetized black hole may behave as a particle accelerators. The BSW process was considered for regular black hole, BTZ black hole and Einstein-Maxwell-Dilation-Axion  in \\cite{huss}. The effect of particle accelerations and collisions on the near horizon surface of some black holes was discussed in \\cite{sharif}.  Also McWilliams \\cite{mc} showed that the black holes are neither particle accelerators nor dark matter probes.  Galajinsky \\cite{gala} also showed that the CM  energy in the context of the near horizon geometry of the extremal Kerr black holes and proved that the CM energy is finite for any value of the particle parameters. Tursunov et al. \\cite{tur} studied the particle accelerations and collisions in case of black string. Fernando \\cite{fernando} has studied the possibility of high CM energy of two particles colliding near the infinite red-shift surface of a charged black hole in string theory. Studies of BSW effect were performed in \\cite{patil5,patil7} for the naked singularity case of different space-times.  In a recent work \\cite{pp3}, we  have demonstrated that an extremal static, spherically symmetric string black hole could act as a particle accelerator with arbitrarily high energy when two uncharged particles falling freely from rest to infinity on the near horizon. It was also shown there that, the CM energy of collision is independent of the extreme  fine tuning of the angular momentum parameter of the colliding particles.  In a present work, our objective is to show an analogous effect of particle collisions with a high CM energy is also possible when the black hole is precisely \\emph{extremal regular} static and spherically symmetric black hole.   Traditional black holes like Schwarzschild black hole, RN black hole, Kerr black hole and Kerr-Newman black holes have possessed a curvature singularity. Whereas the regular black hole does not possess any curvature singularity.  This is the main difference between the traditional black hole and the regular black hole. On the other hand, it is also important to understand the final state of gravitational collapse of initially regular configurations we need to study the global regularity of black hole solutions.  Broadly speaking, for traditional black holes the curvature invariants $R, R_{ab}R^{ab}, R_{abcd}R^{abcd}$ blows up at $r=0$ in the spacetime manifold, while for regular black holes the curvature invariants do not blow up everywhere in the spacetime manifold including at $r=0$. In this sense it is said ``regular black hole'' or ``non-singular black hole''.  We have considered here a number of regular black holes like: the Bardeen black hole \\cite{bard}, Ay\\'{o}n-Beato \\& Garc\\'{i}a (ABG) black hole \\cite{abg1} and Hayward black hole \\cite{hay}.  The special features of these black holes are they have satisfied the weak energy condition(WEC) and the energy-momentum tensor should have the symmetry  $T_{00}=T_{11}$. Our goal here is to see what happens the BSW effect for these black holes in the \\emph{extremal limit}.  Circular geodesic motion in the equatorial plane is of fundamental interest in black hole physics as well as in accretion disk physics to determine the important features of the spacetime. Circular orbits with $r>r_{ISCO}$ are stable, while those with $r<r_{ISCO}$ are not. Keplerian circular orbits exist in the region $r>r_{ph}$, with $r_{ph}$ being the circular photon orbit. Bound circular orbits exist in the region $r>r_{mb}$, with $r_{mb}$ being the marginally bound circular orbit, and stable circular orbits exist for $r>r_{ISCO}$, with $r_{ISCO}$. We have calculated these orbits for the above regular black holes.  These orbits are very crucial in black hole physics as well as in astrophysics because they determine important information on the back ground geometry.  There are a number of references we would like to mention here which discuss the some interesting properties of the regular black holes. Firstly, Ansoldi \\cite{ansol} gives a good review about the regular black hole. Balart \\cite{lb} studied the Brown-York quasi-local energy and  Komar energy at the horizon which is called Bose-Dadhich identity \\cite{bose}, and proved that this identity does not satisfy for regular black holes. In \\cite{nora}, the author showed that Smarr's formula do not satisfy in case of the regular Bardeen and ABG space-time. In \\cite{bala}, the authors have studied the weak energy condition of the regular black hole. Recently, Garc\\'{i}a et al. \\cite{eva} have discussed the complete geodesic structure of ABG spacetime by using Weierstrass elliptic functions. Finally, Eiroa et al. \\cite{eiro} have discussed the gravitational lensing effect of the regular Bardeen black hole.  The paper is organized as follows. In section 2, we compute the CM energy of the particle collision near the infinite red-shift surface of the Bardeen Space-time. In section 3, we describe the BSW effect of the ABG space-time. In section 4, we calculate the $E_{cm}$ for the Hayward black hole and finally we summarize the results in section 5. ",
        "conclusions": "In this work, we have demonstrated that the collision of two neutral particles falling freely from rest at infinity in the background  of the regular black holes. Firstly, we have examined the particle acceleration and collision near the infinite red-shift surface of the Bardeen space-time which is the \\emph{first} regular (singularity-free) black hole model in GR. We proved  that the center of mass energy is arbitrarily large at the infinite red-shift surface when the black hole is purely extremal. For non-extremal Bardeen black hole, we found that the CM energy is finite and depends upon the critical values of the angular momentum parameter.  Secondly, we have examined the BSW mechanism near the infinite red-shift surface of the ABG black hole which is also  a regular black hole space-time and singularity free solutions of the coupled system of a non-linear electrodynamics and general relativity. We have showed that the center of mass energy for this black hole also is arbitrarily large at the infinite red-shift surface when it is in the purely extremal situation. For non-extremal ABG black hole, we have seen that the CM energy is finite and depends upon the fine tuning condition of the angular momentum parameter.  Finally, we have tested the BSW mechanism near the infinite red-shift surface of the Hayward black hole which is also a singularity free solution in GR. It is shown that these regular black holes may act as natural particle accelerators with arbitrarily high CM energy when the black hole is precisely extremal.  Moreover, in each cases, we have also studied the properties of equatorial circular geodesic motion by extremization of  the effective potential for time-like circular orbits and null circular orbits. We particularly emphasized on the  ISCO, MBCO and CPO of these regular black holes. These orbits are useful to extract the information about the back ground geometry and they are also relevant to the astrophysical process.  Our conclusion is that for non-extremal regular space-time  the CM energy is finite and depends upon  the angular momentum parameter. For extremal regular black holes, the CM energy is unlimited due to the diverging values of the angular momentum parameter of the colliding particles."
    },
    "1402/1402.5959_arXiv.txt": {
        "abstract": "We present results of new photometric observations of the contact binary system \\astrobj{HI~Pup} as well as the radial velocity curve of the system. Time series multicolour photometry was obtained at the South African Astronomical Observatory (SAAO) using the 1-m Cassegrain Telescope. We applied a simultaneous solution to the $BVRI$ light and radial velocity curves in order to determine the physical parameters of the system. From an analysis of the new multicolour data, the physical parameters were found to be $M_1=1.22M_{\\odot}$, $M_2=0.23M_{\\odot}$, $R_1=1.44R_{\\odot}$, $R_2=0.67R_{\\odot}$, $L_1=3.3L_{\\odot}$, $L_2=0.7L_{\\odot}$. Our solution confirms that HI Pup has a typical A--type W~UMa binary system characteristics. ",
        "introduction": "The variability of \\astrobj{HI~Pup} (=HIP~36762) was first discovered by \\cite {hof49}. \\cite{sah63} classified this target as an eclipsing binary sytem showing a W|UMa type light curve. Furthermore, they reported \\astrobj{HI~Pup} to be a possible member of galactic cluster \\astrobj{Cr~173}, with a period of 0$^d$.4. The first photoelectric light curve of HI Pup  was obtained by \\cite{ker84}. They confirmed W~UMa--type variation in the light curve of the system  and found the orbital period to be 0$^d$.43257. The star was in the field of {\\it HIPPARCOS} satellite \\citep{esa97} \\cite{ruc97} calculated absolute magnitude of the system using the data taken by {\\it HIPPARCOS}. \\cite{ade01} listed the {\\it HIPPARCOS} targets that have largest photometric amplitudes. \\cite{pri03} included \\astrobj{HI~Pup} into their catalogue of the field contact binary stars. \\cite{due07} obtained and fitted the radial velocity curve of the star. The spectroscopic mass ratio of the system was calculated to be 0.19 by the authors. The physical parameters of 62 binaries including \\astrobj{HI~Pup} were calculated by \\cite{deb11} using the ASAS (All Sky Automated Survey, \\cite{poj03}) database. They obtained the light elements and the absolute parameters of the system by using the radial velocity curve of \\cite{due07}.  In this paper, we present the details of our observations and reduction process in the next section. The simultaneous light and radial velocity curve analyses are explained in Sec. 3. We briefly discuss our results in the last section. ",
        "conclusions": "We analyzed the new multicolour light and radial velocity \\citep{due07} curves of \\astrobj{HI~Pup} to derive the physical parameters of the system that are listed in Table~\\ref{tababs}. The location of the components, corresponding to the determined physical parameters, in the HR diagram is shown in the top panel of Fig.~\\ref{fighr} The less massive and cooler component is located below the Zero Age Main Sequence (ZAMS) while the primary component is close to the secondary one.  The degree of contact of the system, $f$, was calculated by following the equation $f=\\frac{\\Omega_{I}-\\Omega}{\\Omega_{I}-\\Omega_{O}}$ where $\\Omega_{I}$ and $\\Omega_{O}$ are inner and outer Lagrangian surface potential values, respectively. The result is found to be $f$=0.2.  \\cite{bin70} reported that the occultation of the massive component during the primary minimum is a characteristic property of the light curves of A--type W~UMa. The author also categorized two sub-classes of W~UMa systems based on the spectral types of the systems: (i) A--type W~UMa systems which lay between A9 and F8 and (ii) W--types which appear in the spectral type range F7--M5. The evolutionary status of these two type of W~UMa systems are remarkably different from each other. Early studies propose that A--type W~UMa systems are more evolved than W--types. However, \\cite{mac96}, \\cite{gaz06} proved that the idea was no longer accepted. \\cite{mac96} suggested three class for contact binaries: young late--type contacts, old late--type contacts and the rest of contact systems. They remarked that the young and old late--type contact binaries can be considered as W~UMa type. However, third class are related to hot contact systems and they have different evolutionary status. \\cite{gaz08} suggested that these two sub-types should probably belong the group of cool contact binaries and  A--type W~UMa system with high mass can preserve its contact configuration during its evolution. Medium mass W--type binaries, on the other hand, evolve to A--type systems with extreme mass ratios.  The massive component is occulted at primary minimum in the case of HI~Pup. Therefore, the system shows a typical A--type W~UMa light curve. The spectral type given by \\cite{due07}, F6, also prove its membership of A--type according to the classification of \\cite{bin70}. We marked the components of \\astrobj{HI~Pup} with the components of other A--type systems \\citep{lie08} on the mass-radius plane. In the bottom panel of Fig.~\\ref{fighr} it can be clearly seen that the primary and secondary components of A--type W~UMa systems gather in two separate locations in mass--radius diagram.  \\begin{table} \\caption{The derived absolute parameters of the system with the analysis. The effective temperature of the sun is adopted to 5780~K during the calculations.} \\label{tababs} \\begin{tabular}{lrr} \\hline Parameter                &\\multicolumn{2}{c}{Value}    \\\\ & \\multicolumn{1}{c}{Pri.}&\\multicolumn{1}{c}{Sec.} \\\\ \\hline M (M$_{\\odot}$)          & 1.21(24)         & 0.23(19) \\\\ R (R$_{\\odot}$)          & 1.44(11)          & 0.67(10)\\\\ L (L$_{\\odot}$)\t\t & 3.3(5)\t     & 0.7(2) \\\\ M$_{bol}$ ($^m$)         & 3.4(9)\t     & 5.2(1.7) \\\\ $a$ (R$_{\\odot}$)        & \\multicolumn{2}{c}{2.7(2)}  \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure} \\includegraphics{MR_HR.eps} \\caption{Location of the components on the Hertzsprung--Russell diagram (top). The components are placed on the mass--radius plane with various A--type W~UMa binaries given by \\cite{lie08} at the bottom panel. Primary components are symbolised by filled circles while the secondaries are represented with open circles. Squares correspond to components of \\astrobj{HI~Pup}. ZAMS data are taken from \\cite{pol95}.} \\label{fighr} \\end{figure}"
    },
    "1402/1402.0874_arXiv.txt": {
        "abstract": "{ Analyses of recent observations of the sole classical Cepheid in the Kepler field, V1154~Cygni, found random changes  of about 30~minutes in the pulsation period.  These period changes challenge standard theories of pulsation and evolution because the period change is non-secular, and explaining this period jitter is necessary for understanding stellar evolution and the role of Cepheids as precise standard candles.  We suggest that convection and convective hot spots can explain the observed period jitter.  Convective hot spots alter the timing of flux maximum and minimum in the Cepheid light curve, hence change the measured pulsation period. We present a model of random hot spots that generate a localized flux excess that perturbs the Cepheid light curve and consequently the pulsation period which is consistent with the observed jitter.  This result demonstrates how important understanding convection is for modeling Cepheid stellar structure and evolution, how convection determines the red edge of the instability strip, and just how sensitive Cepheid light curves are to atmospheric physics. } ",
        "introduction": "Classical Cepheids are arguably archetypical standard candles since the discovery of the Cepheid period-luminosity relation or Leavitt law more than a century ago \\citep{Leavitt1907}.  Because of the brightness of Cepheids, $\\log L/L_\\odot = 3$ -- $4.5$, and the fact that brightness is correlated with the pulsation period, Cepheids are powerful tools for measuring distances to galaxies \\citep[e.g.][]{Gerke2011,Gieren2013} and even the Hubble constant \\citep{Hubble1929, Freedman2001, Riess2011, Freedman2012}. An additional strength of Cepheids as standard candles is their consistent and precisely measured pulsation periods.  Cepheid pulsation periods are also observed to change over time. \\cite{Eddington1919} presented the first measurements of secular period change in the prototype $\\delta$ Cephei to confirm previous arguments that period change is a result of a central energy source that is not gravitational contraction.  It is now known that the rate of period change is a direct measure of stellar evolution \\citep[e.g.][]{Struve1959} and can test the details of Cepheid evolution, especially Cepheid mass loss \\citep{Neilson2012a, Neilson2012b}.  The measured changes of pulsation periods presented by \\cite{Turner2006} are secular and slowly changing over many years and raise many question such as the nature of the evolution of the nearest Cepheid Polaris \\citep{Neilson2012a, Turner2013a, vanLeeuwen2013}.  New observations are complicating this picture of secular evolution and secular period change. \\cite{Derekas2012} presented Kepler observations of V1154~Cygni where they found that the pulsation period varies by about 0.015 - 0.02 day from its standard pulsation period of about 4.9 days, that is, $\\Delta P/P \\le 1\\%$.  This period jitter is unexpected and presents challenges for detecting binary companions as well as for understanding the physics of Cepheids. \\cite{Evans2013} also found similar period jitter from MOST observations of two other Cepheids and that there apparently is more jitter for first-overtone Cepheids. \\cite{Derekas2012} made several suggestions to explain this jitter, including an instability in the pulsation period itself, or possibly convective granulation in Cepheids, analogous to convective cells observed in red supergiant stars like Betelgeuse \\citep{Gray2008,Haubois2009}.  Convection is known to be important for pulsation in Cepheids, because it defines the red edge of the instability strip \\citep[e.g.][]{Yecko1998, Bono1999}.  Convection is a particularly arduous topic to explore in classical Cepheids and other stars as it is intrinsically a multidimensional problem. Typically, Cepheid structures are computed in one-dimension with various treatments of convection and numerous free parameters \\citep[e.g.][]{Buchler2002, Smolec2008}. New efforts are considering multidimensional simulations of Cepheid structures. \\cite{Gastine2011} computed two-dimensional stellar models to explore the interaction between convection and pulsation in a Cepheid.   However, their model also assumes free parameters akin to one-dimensional models. Similarly, \\cite{Geroux2013} presented two-dimensional radiation hydrodynamic models of convection in RR Lyrae stars and found that simulated light curves matched observational light curves. \\cite{Mundprecht2012} computed three-dimensional Cepheid models that hint at the presence of granulation.  While these models are promising, they are also computationally expensive, requiring both high spacial and temporal resolution, even more so than that required for simulating convection in red supergiant stars \\citep{Freytag2013}.  In this work, we construct a model for convective hot spots that vary as a function of time, number, and size, producing a radiant contribution that is superimposed on a Cepheid light curve.  In Sect.~2, we describe the model for convective spots and pulsation.  In Sect.~3, we present results for a fundamental-mode model light curve, and results for a first-overtone model light curve in Sect.~4.  We summarize our results and discuss implications in Sect.~5. ",
        "conclusions": "The purpose of this work was to understand whether hot convective cells and convective variations might cause the previously observed period jitter \\citep{Derekas2012}.  We computed random hot-spot variations given a number of free parameters such as spot size, spot temperature, spot lifetime, and the number of spots. We found flux variations ranging from $< 0.0001$ mag to about $0.002$ mag and shifts in light minimum up to $\\Delta \\phi = 0.01$ for a fundamental-mode Cepheid and $0.02$ for a first-overtone Cepheid.  These results are more extreme than that observed. The Kepler space telescope is able to resolve flux variations of $0.0001$ mag and a period jitter of about $0.005$, lower than the limits we computed.  However, our model is simple and assumes many degenerate free parameters. The model also adopts a Cepheid light curve and then superimposes convective spot fluxes.  In reality, the Cepheid light curve that was constructed from observations \\citep{Pejcha2012} would contain the average convective flux, so variation due to spots would both increase and decrease the relative flux, hence we are arguably overestimating the flux variation.  Similarly, small changes in our free parameters, such as increasing the number of spots, including limb-darkening and even rotation (albeit slow) could also play a role.  The hypothesis that period jitter is caused by convective hot spots implies that period jitter is itself a function of pulsation period.  We computed the period jitter for a 10-day-period Cepheid, but the period was assumed only for the purpose of constructing the initial light curve. In principle, the amount of relative period jitter depends primarily on the size and temperature of the hot spots. Convection is a function of effective temperature and becomes stronger at cooler effective temperatures.  This suggests that hotter Cepheids would have smaller hot spots and a weaker relative temperature increase (i.e., $T_{\\rm{spot}}/T_{\\rm{eff}}$) than cooler Cepheids.  Because there is a linear relation between the effective temperature and pulsation period, as suggested by the Cepheid period-color relation \\citep[e.g.][]{Tammann2003}, there will be more relative jitter at longer periods.  This suggests that long-period Cepheids, such as $l$~Car and RS~Pup, will have more relative period jitter than short-period Cepheids such as $\\delta$~Cep and  V1154~Cyg. Furthermore, this period jitter may complicate distance estimates from ultra-long-period Cepheids ($P>80$~days) \\citep{Bird2009, Fiorentino2012}.  On the other hand, it may be possible to use the observed period jitter to constrain fundamental stellar parameters similar to the method suggested by \\cite{Bastien2014}.  As convection and convective instability is a function of effective temperature and gravity, then so must be the relative amount of jitter.  One can then combine the amount of period jitter with the average period to help calibrate the Cepheid period-luminosity relation, and fundamental properties for constraining stellar evolution models and multidimensional simulations.  Convective hot spots in Cepheids can be observed by various methods such as interferometric and polarization measurements.  More robust interferometric measurements of the third and fourth lobes of the visibility curve of a Cepheid would constrain perturbations in the intensity profile due to star spots, analogous to the observations by \\cite{Haubois2009}.  Similarly, linear polarization measurements will constrain asymmetric structures on the stellar surface \\citep{Schwarz1984, Clarke1984}. The combination of the two observations will help constrain the presence and evolution of convective cells and confirm the toy model presented in this work.  While our model is simple, the results are robust enough to conclude that convection and convective cells explain the observed period jitter.  However, detailed tests require three-dimensional radiation hydrodynamic simulations of Cepheid atmospheres and the interaction between pulsation and convection.  \\cite{Fokin1996} computed one-dimensional RHD models for $\\delta$ Cep and found that numerous shocks propagated throughout the photosphere, reaching velocities of up to three times the sound speed, hence adding another challenge to computing multidimensional models. Even with these challenges, new three-dimensional models of Cepheids are being computed, and this new field is still in its infancy \\citep{Mundprecht2012}.  In the future, it will be possible to compute grids of pulsating Cepheid atmospheres and potentially verify this work."
    },
    "1402/1402.4720_arXiv.txt": {
        "abstract": "{} {The Sun shows strong variability in its magnetic activity, from Grand minima to Grand maxima, but the nature of the variability is not fully understood, mostly because of the insufficient length of the directly observed solar activity records and of uncertainties related to long-term reconstructions. Here we present a new adjustment-free reconstruction of solar activity over three millennia and study its different modes. } {We present a new adjustment-free, physical reconstruction of solar activity over the past three millennia, using the latest verified carbon cycle, $^{14}$C production, and archeomagnetic field models. This great improvement allowed us to study different modes of solar activity at an unprecedented level of details. } {The distribution of solar activity is clearly bi-modal, implying the existence of distinct modes of activity. The main regular activity mode corresponds to moderate activity that varies in a relatively narrow band between sunspot numbers $\\approx 20$ and 67. The existence of a separate Grand minimum mode with reduced solar activity, which cannot be explained by random fluctuations of the regular mode, is confirmed at a high confidence level. The possible existence of a separate Grand maximum mode is also suggested, but the statistics is too low to reach a confident conclusion. } {The Sun is shown to operate in distinct modes - a main general mode, a Grand minimum mode corresponding to an inactive Sun, and a possible Grand maximum mode corresponding to an unusually active Sun. These results provide important constraints for both dynamo models of Sun-like stars and investigations of possible solar influence on Earth's climate. } ",
        "introduction": "Solar variability has been the subject of intense studies for a long time \\citep[e.g][]{hathawayLR}, but the physics behind it is not yet fully understood \\citep{charbonenau10}. A particularly important question is whether the so-called Grand minima and maxima reported in historical records correspond to special states of the solar dynamo or result from random variability \\citep[e.g.,][]{nandy11,choudhuri12,carbonell94}. Direct records of solar activity, provided by sunspot numbers since 1610 AD \\citep{hoyt98}, are too short to fully address this problem. Reconstructing its past history over longer periods is thus crucial. Cosmogenic radionuclides, such as $^{14}$C, provide powerful proxies \\citep{beer12,solanki_Nat_04,usoskin_LR_13} able to extend the record several millennia back in time. However, previous reconstructions based on such proxies required an ad-hoc calibration, that lead to much debate \\citep{muscheler_Nat_05,solanki_Nat_05}. In the present study, we finally overcome this ambiguity and present the first fully adjustment-free physical reconstruction of solar activity, using the latest carbon cycle \\citep{roth13}, $^{14}$C production \\citep{kovaltsov12}, and archeomagnetic field \\citep{licht13} models, and converting $^{14}$C data into a three-millennia-long sunspot number record. We also study the statistics of the solar activity level to search for distinct modes. ",
        "conclusions": "The cause of the strong variations in solar activity, and whether these result from pure random fluctuations, is still a subject of much debate \\citep[e.g.][]{charbonenau10,choudhuri12,sokoloff04,moss08}. While modern dynamo models can reproduce many observed features of the solar cycle, including secular variability and occurrence of seemingly random Grand minimum events \\citep[e.g.][]{brandenburg08,moss08,choudhuri12}, this often requires a number of ad hoc assumptions \\citep{pipin12}. Thanks to the present reconstruction of solar activity, however, a much clearer picture of Grand minimum events now emerges. These events cannot be described in terms of random fluctuations of a single solar-activity mode. Instead, the evidence suggests that they are manifestations of a Grand minimum mode, distinct from the central regular mode, and that they occur as a result of sudden transitions from this central mode. A similar pattern can be suggested for Grand maximum events as well, but the low statistics does not yet make this suggestion robust. Transitions among these modes can then be understood either as transitions between different branches of the solar dynamo, or as sudden changes in the governing parameters of this dynamo, as proposed by, for example, \\citet{moss08}.  The present result implies that  \\noindent $\\bullet$ the Sun spends most of its time in a regular activity mode with a range of cycle-averaged sunspot numbers of between about 20 and 67  \\noindent $\\bullet$ the Grand minimum state clearly is a separate mode of activity, within which the Sun spends about 1/6 of its time  \\noindent $\\bullet$ there is an indication that the Grand maximum also corresponds to a separate mode of activity, but the low statistics does not allow us to firmly conclude on this, yet.  \\vskip 0.2cm \\noindent These observations provide important constraints for solar-activity models, and will help to deepen our understanding of the processes that drive solar and stellar variability."
    },
    "1402/1402.7045_arXiv.txt": {
        "abstract": "We present results from an original observational campaign comprising five epoch optical photopolarimetrical observations of the BL Lac-type AGN OJ287 in the period 2012 November - 2013 April. The data are gathered with the Focal Reducer Rozhen 2 -- FoReRo2 on the 2-m RCC telescope at NAO Rozhen, Bulgaria. We derive photometry and polarization in R-band, as well as position angle (P.A.). There are indications for correlation between polarization and brightness in R-band. Furthermore, observed variation in P.A. corresponds to a rotation of the plane of polarization of 5.80 $deg$ per day. ",
        "introduction": "Blazars are very powerful and extremely variable sources of polarized radiation. At z=0.306, the BL Lac-type object OJ287  is one of the most well-observed blazars. In fact, observations on its light curve date back to year 1891 (Valtonen \\& Ciprini 2012). But OJ287 is one of the most peculiar AGNs, as well. It shows well-defined period of 12 years (Sillanp\\\"{a}\\\"{a} et al. 1996). Different theoretical models has been proposed to explain its characteristically double-peaked flaring activity during outburst (Katz 1997; Villata et al. 1998; Sillanp\\\"{a}\\\"{a} et al. 1988; Sundelius et al. 1997; Valtaoja et al. 2000). The favored current  model is of a binary black hole system with relativistic precession (Valtonen \\& Ciprini 2012). It is interesting to note that, according to Neronov \\& Vovk (2011), the relativistic beamed emission may come from the jet, produced by the smaller secondary black hole.  The analysis of polarization behaviour plays a key role in the study of blazars (Barres de Almeida et al. 2010). Polarization measurements for OJ287, especially during outburst (Valtaoja et al. 2000), are of crucial importance. The motivation is that the first of the two flares in OJ287 during outburst is thermal (i.e. not polarized). This is shown by the lack of corresponding radio emission (Valtaoja et al. 2000). But the second flare is model-dependent and could be due either to synchrotron radiation, which is polarized (Valtaoja et al. 2000), or can be unpolarized or low-polarized bremsstrahlung (Valtonen \\& Ciprini 2012). The polarization is a key factor to distinguish among different theoretical models.  A number of multiwavelength observational campaigns have been performed in the last years (Valtaoja et al. 2000; Efimov et al. 2002; Ciprini et al. 2007; Nieppola et al. 2009; Villforth et al. 2010; Valtonen \\& Ciprini 2012). Results from the 2005-2010 observational campaign show that the flare during the 2005 outburst is due to bremsstrahlung radiation instead of a synchrotron one (Valtonen \\& Ciprini 2012). That provides strong background to the binary black hole model with ultra relativistic precession. This also allows for the use of OJ287 as a test of general relativity (Valtonen \\& Ciprini 2012). Surprisingly, in April 2012 OJ287 showed a peak in brightness (Santangelo 2012) that was not expected by any of the current models. Possible theoretical explanations according to the binary black hole model can be found in Pihajoki et al. (2013).  There are indications that position angle (P.A.) change is between 2 $deg$ (D'arcangelo et al. 2009) and 5 $deg$ per day (Efimov et al. 2002). But there is a real shortage of optical photopolarimetrical campaigns. Possible correlation between optical flux and polarization degree for OJ287 is suggested by observations from Takalo et al. 1994, but is not well estabished so far (Jorstad et al. 2007, Villforth et al. 2010).  Complete data of polarization behaviour and P.A. measurements, combined with simultaneous photometrical data of OJ287, are important for understanding the underlying physics of this peculiar object. But our work is not aimed at discussing the various theoretical models for OJ287. Instead, we focus on original observational photopolarimetrical study and, as a result, we get some exciting new insights. ",
        "conclusions": "The Blazar-type AGN OJ287 is one of the most well-observed blazars. Yet, it still remains a riddle. Numerous models have been proposed to explain its observed properties (Valtonen \\& Ciprini 2012). Since optical polarization measurements could be useful to distinguish among various theoretical models, our observations further complete the existing data on this peculiar object. Our original 5-epoch study of the optical photometry, polarization and P.A. of OJ287 shows good correlation with previous data  and gives strong background to further measurements with NAO Rozhen's focal reducer FoReRo2 on the 2-m RCC telescope. We observe variation in P.A. that corresponds to rotation of the plane of polarization of 5.80 $deg$ per day. Thus, our work builds on and impoves recent previous research on OJ287 (D'arcangelo et al. 2009; Bozhilov et al. 2013). We also observe dependence of the polarization behaviour with change of optical brightness, as expected by the Binary Black Hole model (Valtonen \\& Sillanp\\\"{a}\\\"{a} 2011)."
    },
    "1402/1402.0981_arXiv.txt": {
        "abstract": "{ We present new and accurate optical photometry of the Bulge globular cluster \\ngc6528. The images were collected with ACS at HST and together with WFC3 (UVIS, IR) allowed us to measure the proper motion to separate cluster and field stars. We adopted two empirical calibrators and we found that \\ngc6528 is coeval with and more metal-rich than 47\\,Tuc. Moreover, it appears older and more metal-poor than the super-metal-rich old open cluster \\ngc6791. We also performed a preliminary analysis of field stellar populations located around \\ngc6528 and \\ngc6522 by using ground-based near-infrared photometry collected with SOFI at NTT. The comparison of evolved stellar components (red giant branch, red horizontal branch, red clump stars) indicates that they share similar properties in this region of the Baade's Window. ",
        "introduction": "The Bulge is a complex component of the Galactic spheroid whose stellar content shows a broad range of kinematic and chemical components. Different observational campaigns connected either with microlensing events in the direction of the Galactic center (OGLE, EROS, MACHO, \\citep{Uda94,Zha95,Alc95,Aub93}, or with near-infrared (NIR) integrated surface photometry from COBE/DIRBE observations \\citep{Dwe95}, or with star counts from the 2MASS survey \\citep{Lop05}, have found evidence that the Bulge has a boxy/peanut morphology, with the long axis lying onto the Galactic plane at an angle of $\\phi_0\\simeq 20^\\circ$ from the Sun-Galactic center direction \\citep{Rat07,Ger02}. More recent investigations based on resolved stellar populations have also found evidence of a bar by analyzing the (bimodal) distribution of the Red Clump (RC) stars in the NIR color-magnitude diagram (CMD) of some Bulge regions located approximately along the Galactic minor axis at $b < 5^\\circ$. \\cite{McW10}, \\cite{Nat10} and \\cite{Sai11} found evidence that RC stars are split into two components separated, on average, by 0.65\\,mag along the line of sight. \\cite{McW10} argue that such a separation can be associated with an underlying X-shaped structure of the Bulge. \\par Although the Bulge cannot be considered a single stellar system, the bulk of its stars are mostly old with ages of the order of 10 - 12\\,Gyr \\citep{Cla08,Sah06,Zoc01} thus suggesting that it is the earliest massive component of Galaxy formed over a short time interval \\citep{McW10,Mel08,Zoc03,Ort95}. Recent empirical and theoretical investigations support the working hypothesis that the Bulge, and in particular the inner Bulge, was the first component of the spheroid to be formed and then it served as a nucleus around which the rest of the Galaxy was built. This is the so-called inside-out scenario \\citep{Zoc08,Lee92}. According to this framework, the innermost regions of the stellar halo and of the Bulge were formed by a rapid, dissipative collapse of low-angular momentum material \\citep{Egg62}. A high-rate of star formation, in the first $\\sim$1\\,Gyr, caused a fast chemical enrichment of the interstellar medium mainly by type\\,II SNe ejecta, as confirmed by the overabundance of $\\alpha$-elements measured in Bulge stars \\citep{Zoc06,McW94}. More recently, the large spectroscopic survey of the Bulge \\cite[ARGOS,][]{Fre13} using RC stars as stellar tracers found that the metal-rich stellar component ([Fe/H] $> -0.5$) is associated with the boxy Bulge, while the metal-poor one ([Fe/H] $< -0.5$) is associated with the thick disk. The former component is made of two distinct populations: the metal-poor component ([Fe/H] $\\approx -0.25$) is uniformly distributed across the selected fields while the metal-rich component ([Fe/H] $\\approx 0.15$) kinematically colder and closer to the plane of the Galaxy. The above components are considered the aftermaths of the instability-driven Bulge formation from the pristine Galactic thin disk \\citep{Nes13-1}. \\par The observation of a radial metallicity gradient in the Bulge \\citep{Nes13-2,Utt12,Zoc08} further support the hypothesis that our Bulge is more a boxy/peanut Bulge than a classical merger-generated Bulge \\citep{Nes13-2,Zoc08}. The Bulge GCs clusters play in this context a key role, since they are the relics of the time when the Bulge forming instabilities took place. In the following we briefly discuss the properties of \\ngc6528 and four typical Baade's Window regions located around \\ngc6528 and \\ngc6522. ",
        "conclusions": ""
    },
    "1402/1402.0665_arXiv.txt": {
        "abstract": "We discuss the possibility that a sterile neutrino of mass around $50$ MeV slightly mixed with the muon flavor may be the origin of the MiniBooNE anomaly. We show that its production in the atmosphere in a fraction of kaon decays would imply an excess of contained showers at IceCube from down-going and near-horizontal directions. ",
        "introduction": "During the past 20 years a number of experiments with solar, atmospheric, reactor and baseline neutrinos have shown that neutrinos have masses and mixings \\cite{Beringer:1900zz}. These experiments provide a framework with \\begin{equation} \\left\\{ \\begin{array}{l} \\displaystyle \\Delta m_{12}^2 \\approx 7.6\\times 10^{-5}\\;{\\rm eV^2}\\\\ \\Delta m_{23}^2 \\approx 2.4\\times 10^{-3}\\;{\\rm eV^2}\\\\ \\;\\;\\;\\;\\;\\;\\;\\;\\;\\; \\approx \\Delta m_{13}^2 \\end{array} \\right. \\;\\;\\;\\;\\;\\;\\;\\;\\; \\left\\{ \\begin{array}{l} \\displaystyle \\sin^2\\theta_{12}\\approx 0.30\\\\ \\sin^2\\theta_{23}\\approx 0.50\\\\ \\sin^2\\theta_{13}\\approx 0.025 \\end{array} \\right. \\end{equation} that fits remarkably well most of the data. From a model-building point of view the importance of this discovery cannot be overstated. We should, however, keep in mind that \\begin{itemize} \\item There are some very basic questions with no answer yet. We do not know, in particular, if these masses are purely electroweak (EW) (just like the electron mass) or if neutrinos are Majorana fermions and their mass is revealing a new scale in particle physics (like radiactivity reveals the EW scale): \\begin{equation} y_\\nu \\; H L \\nu^c\\;\\; {\\rm or} \\;\\; {1\\over \\Lambda_\\nu}\\; H H L L\\;? \\end{equation} \\item There has been a \\emph{persistent} anomaly in several experiments with neutrino beams from particle accelerators. Namely, LSND \\cite{Athanassopoulos:1996jb} and MiniBooNE \\cite{AguilarArevalo:2007it,AguilarArevalo:2009xn} have observed an excess of $\\approx 3$ events with an electron in the final state per each 1000 $\\nu_\\mu$ charged-current (CC) interactions. The interpretation of these events in terms of $\\nu_\\mu\\to \\nu_e$ oscillations would force the addition of two sterile neutrinos of mass around 1 eV. Here we will go in a completely different direction: we will introduce a sterile mode, but much heavier and unstable. Let us very briefly review these experimental anomalies and the heavy neutrino proposal. \\end{itemize} ",
        "conclusions": "Neutrino physics has progressed a great deal during the past 20 years, but a few basic questions and some persistent anomalies should still be clarified. We have focused on MiniBooNE and IceCube, that have provided very interesting although still incomplete results. Hopefully, in the near future MicroBooNE and an increased statistics, respectively, will establish the reach of their current results.  We have argued that there is room in neutrino physics for \\emph{serious} departures from the three-flavor picture, in particular, we have discussed a 50 MeV sterile mode that could be a solution to the baseline anomalies.  We find the possible incidence of such neutrino in completely different experiments, like neutrino telescopes, very interesting. The heavy neutrino could be produced in the atmosphere and decay inside the telescope, introducing an excess of contained showers from downgoing and near horizontal directions, some of them contaminated by atmospheric muons. There seems to be room in current IceCube data for this type of events. Telescopes have the advantage that they are basically background free experients: each individual event must be explained in terms of neutrino interactions or atmospheric muons. In this sense, the 50 MeV neutrino is a possibility that that may be worth exploring."
    },
    "1402/1402.7103_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small Transitional disks are objects whose inner disk regions have undergone substantial clearing. The {\\it Spitzer Space Telescope} produced detailed spectral energy distributions (SEDs) of transitional disks that allowed us to infer their radial dust disk structure in some detail, revealing the diversity of this class of disks.  The growing sample of transitional disks also opened up the possibility of demographic studies, which provided unique insights. There now exist (sub)millimeter and infrared images that confirm the presence of large clearings of dust in transitional disks.  In addition, protoplanet candidates have been detected within some of these clearings. Transitional disks are thought to be a strong link to planet formation around young stars and are a key area to study if further progress is to be made on understanding the initial stages of planet formation. Here we provide a review and synthesis of transitional disk observations to date with the aim of providing timely direction to the field, which is about to undergo its next burst of growth as ALMA reaches its full potential. We discuss what we have learned about transitional disks from SEDs, color-color diagrams, and imaging in the (sub)mm and infrared.  We note the limitations of these techniques, particularly with respect to the sizes of the clearings currently detectable, and highlight the need for pairing broad-band SEDs with multi-wavelength images to paint a more detailed picture of transitional disk structure. We review the gas in transitional disks, keeping in mind that future observations with ALMA will give us unprecedented access to gas in disks, and also observed infrared variability pointing to variable transitional disk structure, which may have implications for disks in general. We then distill the observations into constraints for the main disk clearing mechanisms proposed to date (i.e., photoevaporation, grain growth, and companions) and explore how the expected observational signatures from these mechanisms, particularly planet-induced disk clearing, compare to actual observations. Lastly, we discuss future avenues of inquiry to be pursued with ALMA, JWST, and next generation of ground-based telescopes. \\\\~\\\\~\\\\~\\\\~}% ",
        "introduction": "} \\bigskip \\bigskip Disks around young stars are thought to be the sites of planet formation. However, many questions exist concerning how the gas and dust in the disk evolve into a planetary system. Observations of T Tauri stars (TTS) may provide insights into these questions, and a subset of TTS, the ``transitional disks,'' have gained increasing attention in this regard. The unusual SEDs of transitional disks (which feature infrared excess deficits) may indicate that they have developed significant radial structure.  Transitional disk SEDs were first identified by \\citet{strom89,skrutskie90} from near-infrared (NIR) ground-based photometry and IRAS mid-infrared (MIR) photometry. These systems exhibited small NIR and$/$or MIR excesses, but significant mid- and far-IR (FIR) excesses indicating that the {\\it dust} distribution of these disks had an inner hole (i.e., a region that is mostly devoid of small dust grains from a radius R$_{hole}$ down to the central star). They proposed that these disks were in transition from objects with optically thick disks that extend inward to the stellar surface (i.e., Class II objects) to objects where the disk has dissipated (i.e., Class III objects), possibly as a result of some phase of planet formation. A few years later, \\citet{marsh92, marsh93} proposed that such transitional disk SEDs were consistent with the expectations for a disk subject to tidal effects exerted by companions, either stars or planets.  More detailed studies of transitional disks became possible as increasingly sophisticated instruments became available. The spectrographs on board {\\em ISO} were able to study the brightest stars and \\citet{bouwman03} inferred a large hole in the disk of the Herbig Ae$/$Be star HD~100546 based on its SED. Usage of the term ``transitional disk'' gained substantial momentum in the literature after the {\\it Spitzer Space Telescope}'s \\citep{werner04} Infrared Spectrograph \\citep[IRS;][]{houck04} was used to study disks with inner holes \\citep[e.g.,][]{dalessio05,calvet05}. {\\it Spitzer} also detected disks with an annular ``gap'' within the disk as opposed to holes \\citep[e.g.,][]{brown07,espaillat07b}. In this review, we use the term {\\it transitional disk} to refer to an object with an inner disk hole and {\\it pre\\nobreakdashes-transitional disk} to refer to a disk with a gap. For many (pre\\nobreakdashes-)transitional disks, the inward truncation of the outer dust disk has been confirmed, predominantly through (sub)millimeter interferometric imaging \\citep[e.g.,][]{hughes07,brown08,hughes09,brown09,andrews09,isella10b,andrews11}. We note that (sub)mm imaging is not currently capable of distinguishing between a hole or gap in the disk (i.e., it can only detect a generic region of clearing or a ``cavity'' in the disk). Also, it has not yet been confirmed if these clearings detected in the dust disk are present in the gas disk as well.  The combination of these dust cavities with the presence of continuing gas accretion onto the central star is a challenge to theories of disk clearing.  The distinct SEDs of (pre\\nobreakdashes-)transitional disks have lead many researchers to conclude that these objects are being caught in an important phase in disk evolution.  One possibility is that these disks are forming planets given that cleared disk regions are predicted by theoretical planet formation models \\citep[e.g.,][]{paardekooper04, zhu11,dodson11}.  Potentially supporting this, there exist observational reports of protoplanet candidates in (pre\\nobreakdashes-)transitional disks \\citep[e.g., LkCa~15, T~Cha;][]{kraus11,huelamo11}. Stellar companions can also clear the inner disk \\citep{artymowicz94} but many stars harboring (pre\\nobreakdashes-)transitional disks are single stars \\citep{kraus11}.  Even if companions are not responsible for the clearings seen in all (pre\\nobreakdashes-)transitional disks, these objects still have the potential to inform our understanding of how disks dissipate, primarily by providing constraints for disk clearing models involving photoevaporation and grain growth.  In this chapter, we will review the key observational constraints on the dust and gas properties of (pre\\nobreakdashes-)transitional disks and examine these in the context of theoretical disk clearing mechanisms. In \\S 2, we look at SEDs (\\S~2.1) as well as (sub)mm (\\S~2.2) and IR (\\S~2.3) imaging. We also review IR variability in (pre\\nobreakdashes-)transitional disks (\\S~2.4) and gas observations (\\S~2.5). In \\S 3, we turn the observations from \\S 2 into constraints for the main disk clearing mechanisms proposed to date (i.e., photoevaporation, grain growth, and companions) and discuss these mechanisms in light of these constraints. In \\S 4, we examine the demographics of (pre\\nobreakdashes-)transitional disks (i.e., frequencies, timescales, disk masses, accretion rates, stellar properties) in the context of disk clearing and in \\S 5 we conclude with possibilities for future work in this field. \\bigskip \\bigskip \\bigskip ",
        "conclusions": "} \\bigskip  It is thought that practically all stars have planets (see chapter in this volume by {\\it Fischer et al.}). If we start with the reasonable assumption that these planets formed out of disks, then the question is not if there are planets in disks around young stars, but what are their observable signatures. To date we have detected large, optically thin holes and gaps in the dust distribution of disks and theoretical planet clearing mechanisms can account for some of their observed properties. These (pre\\nobreakdashes-)transitional disks have captured the interest of many scientists since they may be a key piece of evidence in answering one of the fundamental questions in astronomy: how do planets form?  While other disk clearing mechanisms (e.g, grain growth, photoevaporation; \\S~3.1) should certainly be at play in most, if not all, disks \\citep[and may even work more effectively in concert with planets;][]{rosotti13}, here we speculate what kinds of planets could be clearing (pre\\nobreakdashes-)transitional disks, beginning with massive giant planets. If a massive giant planet (or multiple planets) is forming in a disk, it will cause a large cavity in the inner disk since it will be very efficient at cutting off the inner disk from the outer disk. This will lead to significant depletion of small and large dust grains in the inner disk and lower accretion onto the central star (\\S~3.2).  This is consistent with (pre-)transitional disks which have submm and NIR scattered light cavities and MIR deficits in the SED (e.g., LkCa~15, GM~Aur, Sz~91, RX J1604-2130; \\S~2.1--2.3). This is also consistent with the lower accretion rates measured for these objects (\\S~4.2). One caveat is that massive planets would be the easiest to detect, yet there has not been a robust detection of a protoplanet in a (pre-)transitional disk yet.  A less massive planet (or fewer planets) would still lead to substantial clearing of the inner disk, but be less efficient at cutting off the inward flow of material from the outer disk (\\S~3.2). This is consistent with those (pre-)transitional disks that have a submm cavity but no MIR deficit in the SED, such as WSB~60, indicating that small dust still makes it through into the inner disk (\\S~2.1--2.3). Interestingly, WSB~60 is in Oph and overall (pre\\nobreakdashes-)transitional disks SEDs are rare in this region. Given that Oph is quite young (age $\\sim 1$\\, Myr), this may be indicative of the effects of dust evolution and$/$or planet growth. Disks in Oph could be in the initial stages of gap opening by planets whereas in older systems, the dust grains in the inner disk have grown to larger sizes and$/$or the planets have grown large enough to more efficiently disrupt the flow of material into the inner disk. The above suggests that most disks with planets go through a pre-transitional disk phase and that there are many disks with small gaps that have escaped detection with currently available techniques.  More theoretical work is needed to explore if all pre-transitional disks eventually enter a transitional disk phase as planets grow more massive and the inner disk is accreted onto the star. Note also that there are some disks which have submm cavities and MIR SED deficits, but no evidence of clearing in NIR scattered light (e.g., SR~21, DoAr~44, RX~J1615$-$3255).  Apparently, there is some important disk physics regarding accretion efficiencies that we are missing. It is necessary to study disks in other star forming regions with ALMA, both to probe a range of ages and to increase the statistical sample size, in order to address questions of this kind.  Detections of protoplanets would be ideal to test the link between planet mass and disk structure.  While much progress has been made understanding the nature of the disks described above, future work may focus more closely on disks with smaller optically thin gaps that could plausibly be created by low-mass planets. The presence of low mass orbiting companions \\citep[$\\sim 0.1 M_J$;][]{paardekooper06} is expected to alter the dust disk significantly with little impact on the gas disk, whereas high mass companions ($\\gtrsim 1 M_J$) can create gaps or holes in the gas disk and possibly alter its dynamics. {\\it Kepler} finds that super-earth mass objects are very common in mature planetary systems (see chapter in this volume by {\\it Fischer et al.}), so it may be that smaller mass objects form fairly commonly in disks and open holes and gaps that are not cleared of dust. These systems are more difficult to identify with SED studies and current interferometers, so potentially there are many more disks with gaps than currently known. High spatial resolution images from ALMA will be able to detect these small gaps, in some cases even those with sizes down to about 3~AU \\citep{wolf05,gonzalez12}.  There are many remaining avenues that researchers can take to fully characterize clearing in disks around TTS. Theoretically, we can simulate the influence of various mass planets on different sized dust particles, which can be compared with observations at various wavelengths to constrain a potential planet's mass. We can use ALMA, VLTI, and JWST to test the full extent of disk holes and gaps as well as their frequency with respect to age. We can also make substantial progress studying the gas distributions in disks in the near future with ALMA, particularly to determine if the structure inferred for the disk from studies of the dust is the same as that in the gas. Gas tracers may also reveal the presence of an orbiting planet via emission from a circumplanetary disk. Lastly, and perhaps most importantly, we need robust detections of protoplanets in disks around young stars. These future advances will help us understand how gas$/$ice giant planets and terrestrial planets form out of disks and hopefully antiquate this review by the time of PPVII.\\\\    \\textbf{ Acknowledgments.} The authors thank the anonymous referee, C. Dullemond, L. Hartmann, and L. Ingleby for insightful comments which helped improve this review.  Work by C.E. was performed in part under contract with Caltech funded by NASA through the Sagan Fellowship Program executed by the NASA Exoplanet Science Institute. Z.Z. acknowledges support by NASA through Hubble Fellowship grant HST-HF-51333.01-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.\\\\  Finally, a special recognition of the contribution of Paola D'Alessio, who passed away in November of 2013. She is greatly missed as a scientist, colleague, and friend.  Among Paola's most important papers are those which addressed the issue of dust growth and mixing in T Tauri disks.  In \\citet{dalessio99}, she showed that disk models with well-mixed dust with properties like that of the diffuse interstellar medium produced disks that were too vertically thick and produced too much infrared excess, while lacking sufficient mm-wave fluxes.  In the second paper of the series \\citep{dalessio01} Paola and collaborators showed that models with power-law size distributions of dust with maximum sizes around 1 mm produced much better agreement with the mm and infrared emission of most T Tauri disks, but failed to exhibit the 10 {\\micron} silicate emission feature usually seen, leading to the conclusion that the large grains must have settled closer to the midplane while leaving a population of small dust suspended in the upper disk atmosphere.  This paper provided the first clear empirical evidence for the expected evolution of dust as a step in growing larger bodies in protoplanetary disks. Another important result was the demonstration that once grain growth proceeds past sizes comparable to the wavelengths of observation, the spectral index of the disk emission is determined only by the size distribution of the dust, not its maximum.  Along with quantitative calculations of dust properties, the results imply that the typical opacities used to estimate dust masses will generally lead to underestimates of the total solid mass present, a point that is frequently forgotten or ignored.  In the third paper of this series \\citep{dalessio06}, Paola and her collaborators developed models which incorporated a thin central layer of large dust along with depleted upper disk layers containing small dust.  The code which developed these models has been used in over 30 papers to compare with observations, especially those from the Infrared Spectrograph (IRS) \\citep{furlan05,mcclure10} as well as the IRAC camera \\citep{allen04} on board the Spitzer Space Telescope, and more recently from PACS on board the Herschel Space Telescope \\citep{mcclure12}, and from the SMA \\citep{qi11}.  Paola's models also played a crucial role in the recognition of transitional and pre-transitional T Tauri disks \\citep{calvet02,dalessio05,espaillat07b,espaillat10}. In many cases, particularly those of the pre-transitional disks, finding evidence for an inner hole or gap from the spectral energy distribution depends upon careful and detailed modeling.  The inference of gaps and holes from the SEDs is now being increasingly confirmed directly by mm- and sub-mm imaging \\citep{hughes09}, showing reasonable agreement in most cases with the hole sizes predicted by the models. Combining imaging with SED modeling in the future will place additional constraints on the properties of dust in protoplanetary disks.  Paola's influence in the community extended well beyond her direct contributions to the literature in over 100 refereed papers.  She provided disk models for many other researchers as well as detailed dust opacities.  The insight provided by her calculations informed many other investigations, as studies of X-ray heating of protoplanetary disk atmospheres \\citep{glassgold04,glassgold07,glassgold09}, of the chemical structure of protoplanetary disks and propagation of high energy radiation \\citep{fogel11,bethell11} and on the photoevaporation of protoplanetary disks \\citep{ercolano08}.  Paola's passing is a great loss to the star formation community and we remain grateful for her contributions to our field."
    },
    "1402/1402.2726_arXiv.txt": {
        "abstract": "Full polarization measurements of 1665 and 1667-MHz OH masers at 261 sites of massive star formation have been made with the Parkes radio telescope.  Here we present the resulting spectra for 157 southern sources, complementing our previously published 104 northerly sources. For most sites, these are the first measurements of linear polarization, with good spectral resolution and complete velocity coverage.  Our spectra exhibit the well-known predominance of highly circularly polarized features, interpreted as $\\sigma$ components of Zeeman patterns.  Focusing on the generally weaker and rarer linear polarization, we found three examples of likely full Zeeman triplets (a linearly polarized $\\pi$ component, straddled in velocity by $\\sigma$ components), adding to the solitary example previously reported.  We also identify 40 examples of likely isolated $\\pi$ components, contradicting past beliefs that $\\pi$ components might be extremely rare.  These were recognised at 20 sites where a feature with high linear polarization on one transition is accompanied on the other transition by a matching feature, at the same velocity and also with significant linear polarization.  Large velocity ranges are rare, but we find eight exceeding 25 \\kms, some of them indicating high velocity blue-shifted outflows. Variability was investigated on timescales of one year and over several decades.  More than 20 sites (of 200) show high variability (intensity changes by factors of four or more) in some prominent features.  Highly stable sites are extremely rare. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.0723_arXiv.txt": {
        "abstract": "{} {We aim at studying the theoretical properties of the regular spacings found in the oscillation spectra of  \\dss. } {We performed a multi-variable analysis covering  a wide range of stellar structure and seismic properties and model parameters representative of intermediate-mass, main sequence stars. The work-flow is entirely done using a new Virtual Observatory tool: \\tool\\ (the VO gateway for asteroseismic models), which is presented in this paper. } {A linear relation between the large separation and the mean density is predicted to be found in the low frequency frequency domain (i.e. radial orders spanning from 1 to 8, approximately) of the main-sequence, delta Scuti  stars' oscillation spectrum. We found that such a linear behavior stands whatever the mass, metallicity, mixing length, and overshooting parameters considered in this work. The intrinsic error of the method is discussed. This includes the uncertainty in the large separation determination and the role of rotation. The validity of the relation found is only guaranteed for stars rotating up to 40\\% of their break-up velocity. Finally, we applied the diagnostic method presented in this work to five stars for which regular patterns have been found. Our estimates for the mean density and the frequency of the fundamental radial mode match with those given in the literature within a 20\\% of deviation.} {Asteroseismology has thus revealed an independent direct measure of the average density of \\dss, analogous to that of the Sun. This places tight constraints on the mode identification and hence on the stellar internal structure and dynamics, and allows a determination the radius of planets orbiting around \\dss\\ with unprecedented precision. This opens the way for studying the evolution of regular patterns in pulsating stars, and its relation with stellar structure and evolution.} ",
        "introduction": "}  Nowadays, thanks to the asteroseismic space missions like \\most\\ \\citep{most03}, \\corot\\ \\citep{Corot03}, and \\kepler\\ \\citep{Gilliland10}, the study of the intrinsic variability of A- and F-type stars is living a revolution. In particular, the large number of modes detected and the variety of frequency domain covered has given rise to the so-called \\emph{hybrid} phenomenon, which imposes a revision of the current observational instability strips of \\ds\\ and \\gd\\ stars \\citep[see ][ and references therein]{Uytt11}. Those stars have been found to be located over the entire \\gd\\ and \\ds\\ instability strips, which implies that a review of their pulsation mechanisms is necessary to supplement (or even substitute) the convective blocking effect and the $\\kappa$ mechanisms, respectively.  From the theoretical side, based on stellar energy balance studies, \\citet{MoyaCris10} concluded that \\dss\\ are able to excite hundreds of pulsation modes, and whose accumulated pulsation energy is not large enough to destroy their hydrostatic equilibrium. Likewise, the richness found in the oscillation spectrum of these stars is interpreted by \\citet{Kallinger10} as non-radial pulsation superimposed on granulation noise (correlated noise).  In this work we focus on \\dss\\ which are intermediate-mass (i.e., from 1.4 to 3$\\,\\msun$, approximately), mainly in the main sequence and the sub-giant branch, but also in the pre-main sequence. Their spectral types range from A2 to F5, and their luminosity classes go from IV to V, approximately. This locates them at the lower part of the Cepheid's instability strip. Their pulsations are mainly driven by the so-called $\\kappa$ mechanism \\citep[see reviews by ][]{Breger00rev, Handler09rev}, showing radial and non-radial oscillation modes excited in a frequency domain ranging from a few tens of $\\muHz$ up to about $800-900\\,\\muHz$ \\citep[see e.g. ][]{Poretti11}.  The interpretation of the oscillation spectra  of these stars has never been an easy task. Due to the complexity of their oscillation spectra the identification of detected modes is often difficult. A unique mode identification is often impossible and this hampers the seismology studies for these stars. Additional uncertainties arise from the effect of rapid rotation, both directly on the hydrostatic balance in the star and, perhaps more importantly, through mixing caused by circulation or instabilities induced by rotation \\citep[see e.g. ][]{Zahn92}.  However some decades ago, this scenario started to change thanks to the detection of regular patterns in the detected frequencies. These patterns were observed in the low radial order modes (mixed modes) frequency domain, in which main-sequence classical pulsators show the maximum of their oscillation power,  i.e. between 80 and 800 $\\muHz$, hereafter intermediate frequency domain. There, regularities similar to those found in solar-like stars were not expected \\citep[for a review on this topic, see][]{Goupil05}. Thanks to great efforts made in long ground-based multi-site campaigns, \\cite{Handler97} found regularities (near $26\\,\\muHz$) in the oscillation spectrum of the young \\ds\\ star \\object{CD-24 7599}, which contained 13 frequencies. Those regularities were found using a Fourier transform technique. Also \\cite{Breger99} found a regular spacing close to $46\\,\\muHz$ in the oscillation spectrum of the \\ds\\ star \\object{FG\\,Vir}, composed of 24 frequencies, searching for  regular spacings using histograms of frequency differences, instead of Fourier transform. Later on, they re-studied the star using data spanned in ten years \\citep{Breger09}, in which the number of detected frequencies increased to 68, approximately. The frequency spacing found was attributed to an observed clustering of certain non-radial modes around radial ones. This clustering was explained by the higher probability of certain $\\ell=1$ modes \\citep[the so-called trapped modes ][]{Dziembowski90}  to be excited to observable amplitudes than other modes.  Thanks to the different space missions providing asteroseismic data, similar studies were undertaken, with much more frequencies detected with very high precision. In 2007, \\citet{Matt07} found 88 frequencies in the oscillation spectrum of the \\ds\\ star \\object{HD209775} observed by \\most. In that work, a regular spacing of $50\\,\\muHz$ approximately was obtained using histograms of frequency differences. This was assumed to be a large separation, although no physical explanation to support this was given. Recently, regular patterns were also found in the oscillation spectra of  \\dss\\ observed by \\corot\\  \\citep{GH09,Manteg12, GH13aa} and \\kepler\\ \\citep{GH13sp}, and the hypothesis of identifying them with the large separation has been reinforced.  Theoretically large separation is expected to grow from low to high radial orders. In contrast to the well-defined plateau shown by the large separation in the high radial order domain \\citep[the so-called asymptotic regime, see e.g. ][]{Antoci11, Zwintz11}, it forms a quasi-periodic structure at low radial order regime \\citep[see Fig.~4 of ][]{GH09}. In that work, it was shown that the standard deviation of such an structure is roughly $2.5\\,\\muHz$, and the distance between the mean value of this structure and the high order regime  is of the order of $5\\,\\muHz$.  The main questions we intend to answer in the present work are: What are the physical properties of the periodicities observed in the low frequency domain of \\dss? Do they have similar properties to those found in solar-like stars?  The above questions are tackled by analyzing the properties of the predicted regularities over a large collection of models and physical variables. To do so, we developed a virtual observatory tool, \\tool\\footnote{http://svo.cab.inta-csic.es/theory/sisms3/index.php?}, designed to easily handle stellar and seismic models, examine their properties, compare them with observational data and find models representative of the studied star(s). The tool is presented in this paper. ",
        "conclusions": "}  We have studied the theoretical properties of the regular spacings in the oscillation spectra of \\dss. A comprehensive dataset of models representative of A-F stars covering a wide range of stellar physical magnitudes (e.g. effective temperature, gravity, metallicity, etc.)  and model parameters has been performed.  We have analyzed the behavior of the predicted large spacings (calculated in the intermediate frequency domain, i.e. the frequency domain in which \\dss\\ pulsate) and their possible dependencies upon other physical magnitudes and/or model parameters. The work-flow has entirely been done using \\tool, a virtual observatory tool for managing and analyzing asteroseismic models that we have developed within the Spanish Virtual Observatory (see Appendix~\\ref{apsec:tool}).  Firstly, we have shown that,  for the frequency domain where \\dss\\ pulsate, the regular spacings obtained using the FT technique (as in GH09) and the regular mathematical definition of large separation computed by \\tool, are only equivalent when selecting oscillation modes with radial order $n\\ge1$.  Secondly, we have found a linear relation between large spacing and the mean density (in a logarithmic scale), for models in the main sequence, whatever the metallicity, mass or convection parameters ($\\amlt$ and $\\dov$) considered (within the values typically used for \\dss).  With the use of a large model grid, we have constructed a diagnostic diagram whose properties have also been studied in detail. In particular, the most important source of error in the determination of the mean density, and other quantities like the frequency of the fundamental radial mode, $\\nu_0$, comes from the uncertainty in the large spacing measure for rotating stars. In particular our results are valid only for stars rotating at most at 40\\% of the break-up velocity. Nevertheless, these uncertainties are expected to be drastically lowered by the inclusion of additional observational constraints (as shown by GH09).  Finally we have applied the diagnostic method presented in this work to five stars for which regular patterns have been found. Our estimates for the mean density and the frequency of the fundamental radial mode match reasonably (within 20\\% of deviation) with those given in the literature using asteroseismology. The comparison with other independent (non-asteroseismic) methods to obtain the mean density of stars reveals that asteroseismology provides, by now, the more precise estimates.   The present work implies a significant step to approach the asteroseismic studies of A-F, pulsating stars, to the level of precision achieved in solar-like stars. This is particularly relevant, not only to understand the structure and evolution of these stars but also for the study of planetary systems. Moreover, the observation of solar-like oscillations in star hosting planets provides a significant help for the characterization of the planetary systems, through the precise knowledge of the mass, radius, mean density and age of the host star \\citep[e.g. ][]{JCD10}. The present work permits the extent of those capabilities to the study of planets orbiting around A-F pulsating stars. In fact, most of the  planets discovered  by direct imaging are found to be orbiting such stars. These systems are critical to understand the spin-orbital interactions between the planets and the hosting star \\citep{Winn09,Wright11}. For these, an estimate of the age of the system is crucial to determine the mass of the planets discovered. Thanks to the large amount of A-F stars observed by CoRoT and Kepler, the number of applications of the relations here derived open a new and important window to study the structure and evolution of these stars \\citep{Uytt11,Moya10}.  In a forthcoming study, we examine the statistical properties of the frequency spectra by analysing their autocorrelation functions and by looking at the cumulative distribution functions of the frequency separations. Results on mode visibility \\citep{Reese13} together with the use of multi-colour photometry will help us to better asses how large separation is affected by rotation, and to improve the accuracy of the relation found in this paper.\\\\[0.1cm]"
    },
    "1402/1402.7335_arXiv.txt": {
        "abstract": "\\noindent Recently, indications for an emission line at $3.55\\,{\\rm keV}$ have been found in the combined spectra of a large number of galaxy clusters and also in Andromeda. This line could not be identified with any known spectral line. It is tempting to speculate that it has its origin in the decay of a particle contributing all or part of the dark matter. In this note we want to point out that axion-like particles being all or part of the dark matter are an ideal candidate to produce such a feature. More importantly the parameter values necessary are quite feasible in extensions of the Standard Model based on string theory and could be linked up to a variety of other intriguing phenomena, which also potentially allow for new tests of this speculation. ",
        "introduction": "Two groups~\\cite{Bulbul:2014sua,Boyarsky:2014jta} have recently found a new emission line in the spectra of galaxy clusters and also in a single galaxy, Andromeda~\\cite{Boyarsky:2014jta}. This line is at an energy of $3.55\\,\\,{\\rm keV}$ and so far has not been identified with a transition in any known element. Although the statistical significance is at the $(4-5)\\sigma$-level, the line is very weak and a number of systematic effects could affect the interpretation. Keeping this caveat in mind, it is nevertheless tempting to speculate that this line could be a sign of new physics, more precisely the sign of a decaying form of dark matter. Indeed, the authors of~\\cite{Boyarsky:2014jta} have found indications that the line becomes stronger towards the center of the cluster, which is consistent with the idea of a decaying dark matter particle.  Both~\\cite{Bulbul:2014sua,Boyarsky:2014jta} as well as~\\cite{Ishida:2014dlp} provide a possible explanation in the form a sterile neutrino decaying into a photon and an active neutrino. Another possibility could be eXciting Dark Matter~\\cite{Finkbeiner:2014sja}. In this brief note we point out, similar to~\\cite{Higaki:2014zua}, that the decay of an axion-like particle (ALP) being all or part of the dark matter is another simple explanation. The strength of the coupling to two photons required can be easily motivated in extensions of the Standard Model based on string theory.  String theory often also features a number of additional ALPs, which generically have similar couplings. Thus, one can link the 3.55 keV line to a variety of other puzzling observations. While this may increase the level of speculation it provides additional handles to test the hypothesis. ",
        "conclusions": "An axion-like particle (ALP) making up all or part of dark matter provides a simple explanation for the recently observed 3.55 keV line in the spectra of stacked galaxy clusters. The required values for the mass and coupling can be obtained in models of string theory. The latter also provide for the possibility that there also is an axion and/or additional ALPs with roughly the same coupling to photons. Using this one can find links to other puzzling astrophysical phenomena and perhaps more importantly also to interesting experimental probes of this hypothesis."
    },
    "1402/1402.6333_arXiv.txt": {
        "abstract": "We investigate lithium depletion in standard stellar models (SSMs) and main sequence (MS) open clusters, and explore the origin of the Li dispersion in young, cool stars of equal mass, age and composition. We first demonstrate that SSMs accurately predict the Li abundances of solar analogs at the zero-age main sequence (ZAMS) within theoretical uncertainties. We then measure the rate of MS Li depletion by removing the [Fe/H]-dependent ZAMS Li pattern from three well-studied clusters, and comparing the detrended data. MS depletion is found to be mass dependent, in the sense of more depletion at low mass.  A dispersion in Li abundance at fixed \\teff\\ is nearly universal, and sets in by $\\sim$200 Myr. We discuss mass and age dispersion trends, and the pattern is mixed. We argue that metallicity impacts the ZAMS Li pattern, in agreement with theoretical expectations but contrary to the findings of some previous studies, and suggest Li as a test of cluster metallicity. Finally, we argue that a radius dispersion in stars of fixed mass and age, during the epoch of pre-MS Li destruction, is responsible for the spread in Li abundances and the correlation between rotation and Li in young cool stars, most well known in the Pleiades. We calculate stellar models, inflated to match observed radius anomalies in magnetically active systems, and the resulting range of Li abundances reproduces the observed patterns of young clusters. We discuss ramifications for pre-MS evolutionary tracks and age measurements of young clusters, and suggest an observational test. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.1039_arXiv.txt": {
        "abstract": "We report \\CSss, \\HCNft, and \\HCOPft\\ observations in 20 nearby star-forming galaxies with the Acatama Pathfinder EXperiment 12-m telescope. Combined with 4 CS, 4 HCN, and 3 HCO$^+$ detections in literature, we probe the empirical link between the luminosity of molecular gas (\\Lgas) and that of infrared emission (\\LIR), up to the highest gas densities (10$^6$ - 10$^8$\\cmt) that have been probed so far. For nearby galaxies with large radii, we measure the IR luminosity within the submm beam-size (14 --18$''$) to match the molecular emission. We find linear slopes for \\LCSss-\\LIR\\ and \\LHCNft-\\LIR, and a slightly super-linear slope for \\LHCOPft-\\LIR. The correlation of \\LCSss-\\LIR\\ even extends over eight orders of luminosity magnitude down to Galactic dense cores, with a fit of log(\\LIR)$=1.00(\\pm 0.01) \\times$log(\\LCSss) $+4.03(\\pm0.04)$. Such linear correlations appear to hold for all densities $>10^4$\\cmt, and indicate that star formation rate is not related to free-fall time scale for dense molecular gas. ",
        "introduction": "The star formation process constantly turns gas into stars. The Kennicutt-Schmidt (K-S) law \\citep{Kennicutt98} globally correlates the surface-densities of star formation rate ($\\Sigma_{\\rm SFR}$ traced by H$\\alpha$) with total gas mass ($\\Sigma$gas traced by CO and \\HI), with a slope $\\alpha \\sim 1.4\\, (\\Sigma_{\\rm SFR} \\propto \\Sigma_{\\rm gas}^\\alpha)$. However, quite a range of deviations from the K-S law have been observed, and no unique slope was found in the \\LCO-\\LIR\\ correlations as well \\citep[e.g.,][]{Liu2012, gs04b}.   Recent studies on star formation indicate that stars, especially the massive stars, are predominantly formed in the dense cores of giant molecular clouds \\citep[e.g.,][]{e08}. Dense gas directly represents molecular content involved in forming stars \\citep[e.g.,][]{Lada12}, traced by the rotational transitions of high dipole moment molecules (e.g., HCN and \\HCOP), because of their high critical densities\\footnote{All critical densities ($n_{\\rm crit}$) in this letter are calculated with $n_{\\rm crit} = \\sum\\limits_{u>l} A_{\\rm u l} / \\sum\\limits_{\\rm u \\neq l} C_{\\rm ul}(T_{\\rm kin})$ at \\Tkin=100 K assuming optically thin emission. $A_{\\rm ul}$ and $C_{\\rm ul}$ denote the Einstein coefficient for spontaneous emission and the collision rate, respectively. All state-to-state cross sections and rate coefficients are from the LAMDA Web site (\\url{http://home.strw.leidenuniv.nl/\\~moldata/}) \\citep{schoier2005}.}(\\ncrit).  \\citet[][]{gs04b,gs04a} find a tight linear correlation between the luminosities of IR emission (\\LIR\\, tracing the SFR) and HCN $J$=1$\\rightarrow$0 (\\LHCNoz\\ tracing $M_{\\rm dense}$) in galaxies. This correlation even extends to Galactic dense cores undergoing high mass star-formation \\citep[e.g.,][]{Wu05}.   \\citet{Lada12} argue that such linear correlation is a fundamental relation for SFR and dense gas, and the molecular gas with densities above $10^4$\\cmt\\ should follow this linearity. Linear correlations have also been found in other dense gas tracers with similar or higher \\ncrit\\ (e.g., \\HCOP $J$=1$\\rightarrow$0, HNC $J$=3\\to2, \\CSff), in both Galactic dense cores and galaxies \\citep[e.g.,][]{Schenck11,Ma12,Reiter2011, Wang2011}.   \\begin{figure*}[btp] \\includegraphics[angle=-180,scale=1.0,bb=1 515 510 790 ]{fig1.ps}   \\caption{Spectra of the strong detections of \\HCOPft. Filled lines show \\CSss, red thick lines show \\HCOPft, and blue dotted lines show \\HCNft. All velocities are labeled relatively to the velocity corresponding to their redshifts. } \\label{spectra1} \\end{figure*}       \\begin{figure*}[btp] \\includegraphics[angle= 0,scale=1.0,bb=1 1 510 280]{fig2.ps} \\caption{Spectra of the weak detections of \\HCOPft. The labels are the same as the Fig \\ref{spectra1}.} \\label{spectra2} \\end{figure*}     \\begin{table*} \\caption{ Observed properties of the sample}\\label{lineinfo} \\begin{center} \\scriptsize \\begin{tabular}{lcccccccccccc}\\hline\\hline &Distance& \\LIR\\ & $I_{\\rm CS(7-6)}$ &$I_{\\rm HCN(4-3)}$ &$I_{\\rm HCO^+(4-3)}$ & $L'_{\\rm CS(7-6)}$& $L'_{\\rm HCN(4-3)}$ & $L'_{\\rm HCO^+(4-3)}$ & $R_{\\rm SD}$&$C_{\\rm aper}$&Band&Type$^g$ \\\\ Source Name & (Mpc)& ($10^{10}$\\Lsun)  & (\\Kkms)  & (\\Kkms)  & (\\Kkms)  &  (10$^6$\\Kkmspc)   &  (10$^6$\\Kkmspc)    &  (10$^6$\\Kkmspc)    &   & &($\\mu$m) &   \\\\ \\hline IC~1623         & 82.4 & 49.0  &    $<$  0.9 &         $<$  0.9 & 0.9$\\pm$0.2 &    $<$  66   &     $<$  62     &    62$\\pm$15   & 0.65 & 1.29 &100&SF\\\\ NGC~1068        & 14.4 & 20.4  & 1.0$\\pm$0.2 &      5.0$\\pm$0.2 & 2.7$\\pm$0.2 & 2.5$\\pm$0.5  & 10.7$\\pm$0.4    &   5.8$\\pm$0.4  & 0.16 & 1.16 &70 &AGN\\\\ NGC~1365        & 18.3 & 10.5  & 0.6$\\pm$0.2 &      1.7$\\pm$0.2 & 2.9$\\pm$0.2 & 2.2$\\pm$0.7  &  5.9$\\pm$0.7    &  10.0$\\pm$0.7  & 0.25 & 1.16 &70 &AGN\\\\ NGC~1808        & 12.3 & 6.92  &    $<$  0.6 &      2.5$\\pm$0.2 & 2.2$\\pm$0.2 &    $<$  1.0  &  3.9$\\pm$0.3    &   3.4$\\pm$0.3  & 0.46 & 1.29 &100&AGN\\\\ NGC~2903        & 10.6 & 2.57  &    $<$  0.9 &         $<$  0.9 &    $<$  0.9 &    $<$  1.1  &     $<$  1.0    &      $<$  1.0  & 1.0  & --   &-- &SF\\\\ NGC~3256        & 44.3 & 57.5  & 0.8$\\pm$0.2 &      1.1$\\pm$0.2 & 2.8$\\pm$0.2 &  45$\\pm$6    &   22$\\pm$4      &    56$\\pm$4    & 0.34 & 1.29 &100&SF\\\\ NGC~3627        & 11.0 & 2.88  &    $<$  0.7 &      0.8$\\pm$0.2 & 0.7$\\pm$0.2 &    $<$  1.0  &  1.0$\\pm$0.3    &   0.9$\\pm$0.3  & 0.08 & 1.29 &100&AGN\\\\ NGC~3628        & 11.0 & 2.14  & 0.5$\\pm$0.2 &      0.9$\\pm$0.2 & 2.5$\\pm$0.2 &    $<$  1.2  &  1.1$\\pm$0.4    &   3.1$\\pm$0.4  & 0.23 & --   &850&SF\\\\ NGC~4039        & 25.0 & 9.33  &    $<$  1.2 &         $<$  0.7 & 0.9$\\pm$0.3 &    $<$  8.2  &     $<$  4.5    &   5.8$\\pm$1.8  & 0.04 & 1.29 &100 &SF\\\\ NGC~4418        & 29.5 & 9.33  & 1.2$\\pm$0.2 &      3.0$\\pm$0.2 & 1.2$\\pm$0.2 &11.5$\\pm$2.0  & 26.8$\\pm$2.0    &   8.9$\\pm$1.8  & 1.0  & --   & --&AGN\\\\ NGC~4945        & 3.8  & 2.82  & 7.3$\\pm$0.3 &     23.4$\\pm$0.3 & 9.9$\\pm$0.3 & 1.2$\\pm$0.05 &  3.6$\\pm$0.05   &   4.5$\\pm$0.05 & 0.40 & 1.29 &100&AGN\\\\ IRAS13120-5453  & 128.3& 178   &    $<$  0.9 &      2.0$\\pm$0.3 & 1.6$\\pm$0.3 &    $<$  106  &  330$\\pm$30     &   265$\\pm$30   & 1.0  & --   &-- &SF\\\\ IRAS13242-5713  & 38.9 & 31.6  & 1.8$\\pm$0.3 &      4.3$\\pm$0.3 & 3.5$\\pm$0.3 &  30$\\pm$5    &   67$\\pm$5      &    54$\\pm$5    & 0.64 & 1.29 &100&SF\\\\ NGC~6240        & 107  & 70.8  &    $<$  0.7 &      1.3$\\pm$0.2 & 2.4$\\pm$0.2 &    $<$  99   &  189$\\pm$30     &   260$\\pm$20   & 1.0  & --   &-- &SF\\\\ IRAS17578-0400  & 61.1 & 24.5  &    $<$  0.9 &      1.5$\\pm$0.3 &    $<$  0.9 &    $<$  36.6 & 57.0$\\pm$10     &      $<$  34   & 1.0  & --   &-- &SF\\\\ IRAS18293-3413  & 77.4 & 64.6  &    $<$  0.9 &      $<$ 0.9     & 1.6$\\pm$0.3 &    $<$  58.6 &     $<$  55     &   116$\\pm$20   & 1.0  & --   &-- &SF\\\\ NGC~7469        & 71.6 & 46.8  &    $<$  0.9 &         $<$  0.9 & 1.0$\\pm$0.2 &    $<$  33.5 &     $<$  30     &    37$\\pm$10   & 1.0  & --   &-- &AGN\\\\ NGC~7552        & 19.5 & 8.91  & 0.6$\\pm$0.2 &      1.3$\\pm$0.2 & 2.4$\\pm$0.2 & 1.3$\\pm$0.4  &  5.1$\\pm$0.4    &   9.4$\\pm$0.4  & 0.59 & 1.29 &100&SF\\\\ NGC~7771        & 63.0 & 26.9  &    $<$  0.9 &         $<$  0.9 &    $<$  0.9 &    $<$  40   &     $<$  36     &      $<$  36   & 0.55 & 1.29 &100 &SF\\\\ Mrk~331         & 80.9 & 33.9  &    $<$  0.9 &      1.1$\\pm$0.2 &    $<$  0.9 &    $<$  43   &   60$\\pm$13     &      $<$  40   & 1.0  & --   &-- &AGN\\\\ \\hline NGC~253$^a$     & 3.47 & 3.45  &     --      &         --       &     --      &0.35$\\pm0.01$$^a$ & 8.9$\\pm$1$^b$   & 7.7$\\pm$1.0$^b$ & 0.18&1.33 &100 &SF\\\\ IC~342$^b$      & 3.93 & 1.08  &     --      &         --       &     --      &0.05$\\pm0.02$$^a$ & 0.8$\\pm$0.2$^c$ &      --         & 0.05&1.33 &100 &SF\\\\ Arp~220$^c$     & 79.9 & 162.2 &     --      &         --       &     --      &190 $\\pm40$$^d$   & 950$\\pm$190     & 170$\\pm$40 $^d$ & 1.0 &--   &--  &SF\\\\ HE2-10$^d$      & 16.4  & 0.6  &      --     &         --       &     --      &0.21$\\pm0.7$$^a$  &    --           &      --         & 0.29&1.33 &100 &SF\\\\ MRK~231$^e$     & 180.8 & 340  &     --      &         --       &     --      &    --            & 550$\\pm$110$^e$ & 188$\\pm$23$^f$  & 1.0 &--   &--  &AGN\\\\ \\hline \\end{tabular}\\\\ $^a$ \\citep[][]{bayet09}. $^b$  \\citep[][]{Knudsen07}. $^c$ \\citep[][]{Jackson95}. $^d$  \\citep[][]{gpg09}. $^e$ \\citep[][]{PPP2007}. $^f$ \\citep[][]{Wilson08} $^g$ Galaxy types are found in NASA/IPAC Extragalactic Database (NED). \\end{center} \\end{table*}   For gas tracers with \\ncrit\\ higher than HCN $J$=1$\\rightarrow$0, the slopes of the \\Lgas-\\LIR\\ correlations are controversial. \\citet{kt07} argue that the mean densities in different types of galaxies and \\ncrit\\ of the tracer change the slopes. Numerical simulations predict decreasing slopes against increasing \\ncrit, because of sub-thermal excitation conditions \\citep[i.e.,][]{ncs08,jnm09}. Observations of HCN, \\HCOP, and CS (typically, $J$$\\le$3$\\rightarrow$2) have been used to support these and show sub-linear correlations \\cite[e.g.,][]{bhl08, bns08,gc08}, where \\HCNtt\\ has a slope of $\\alpha\\sim$ 0.8, following the prediction quite well \\citep{bns08}. Up to date, only few detections of higher-$J$ transitions of the above species have been reported \\citep[i.e.,][]{Knudsen07,Jackson95,gpg09,bayet09, Wilson08}.  To better probe the densest molecular gas in galaxies, and to test the predictions for higher \\ncrit\\ tracers, we therefore performed a survey of \\CSss, \\HCNft, and \\HCOPft\\ in 20 nearby actively star-forming galaxies with the Atacama Pathfinder EXperiment (APEX) 12-m telescope\\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut f{\\\"u}r Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}. With the advent of the {\\it Herschel} space telescope, we are able to obtain beam-matched IR luminosities in nearby galaxies, and to compare them with data from single dish telescopes. In this letter, we summarize our findings and compare them with the results from Galactic studies. We adopt flat cosmology parameters ($\\Omega_m$ =0.27, $\\Omega_\\Lambda$=0.73, and $H_0$=71 \\kms Mpc$^{-3}$; \\citealp{sbd07}). ",
        "conclusions": "We detected 8, 14, and 16 ($I_{\\rm line}$$\\ge$\\ 3 $\\sigma$) galaxies in \\CSss, \\HCNft, and \\HCOPft, respectively. Most of them are strong detections of HCN and \\HCOP\\ ($\\ge$\\ 4 $\\sigma$), while Mrk~331, NGC~7469 and NGC~4039 show marginal detections in \\HCOPft\\ ($\\sim$ 3 $\\sigma$). Upper limits (3 $\\sigma$) are derived for the non-detections. Fig. \\ref{spectra1} shows the spectra of strong detections of \\HCOPft, and Fig. \\ref{spectra2} shows the spectra of weak detections of \\HCOPft. NGC~7771 and NGC~2903 has only non-detections, so we do not show their spectra.  Combining data from the literature we construct samples containing 14, 17, and 17 detected galaxies for \\CSss, \\HCNft, and \\HCOPft, respectively.   Total IR (TIR; 8-1000$\\mu$m) luminosity \\citep{smk03} is adopted as a proxy of SFR.  However, the molecular lines are obtained only in small beam-sizes ($\\sim$ 14$''$ -- 18$''$), which only pick up a certain portion of the IR luminosity for a whole galaxy.  This is particularly problematic for nearby galaxies and starbursts because of their large radii (up to a few arcmins). If this effect was neglected, one would systematically overestimate the IR emission or underestimate the dense gas emission at the low luminosity end of the \\Lgas-\\LIR\\ correlations, while the ULIRGs are not affected due to their small angular sizes of gas emission.  We decide to measure the IR emission within the submm beam-size, rather than to adopt the IR luminosities of the entire galaxies. We download {\\it Herschel}\\footnote{Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.} PACS 100 $\\mu$m, 70 $\\mu$m (when 100 $\\mu$m is not available) images from the Herschel Science Archive (HSA), and they are processed to level 2.5 in the pipeline. NGC 3628 does not have {\\it Herschel} data, so we adopt archival images of SCUBA 850 $\\mu$m instead. Using these IR data, we perform aperture photometry with the submm beam-size and with the whole galaxy. The background radii are selected by eye from the outside of galaxies to the edge of the images. Some nearby galaxies are downloaded from the Key Insights on Nearby Galaxies (KINGFISH) project, and similar to their photometry results \\citep{Dale2012}, the impact of using $\\sim$10\\% larger or smaller aperture areas is a median difference of less than 3\\% on the flux densities for all wavelengths. The IR luminosity within the beam-size is:  \\[L_{\\rm IR} = L_{\\rm TIR} \\times R_{\\rm SD} \\times C_{\\rm aper} ,\\]  where $L_{\\rm TIR}$ is the IR luminosity of the entire galaxy, $R_{\\rm SD}$ is the ratio of flux densities within the beam-size of single dish (SD) telescope to that measured in the whole galaxy, and $C_{\\rm aper}$ is the aperture correction factor for the beam-sizes. The final error comprises errors of photometry ($\\sim$5\\%), the point source assumption ($\\sim$10\\%), the flux calibration error ($\\sim$5\\%), and the error of tracing TIR with a chromatic IR band ($\\sim$10\\%) \\citep[e.g.,][]{Galametz2013}. In the end we account 20\\% as a conservative uncertainty for \\LIR.  \\subsection{The \\Lgas-\\LIR\\ Correlations}  Figure \\ref{correlations} presents the \\Lgas-\\LIR\\, correlations. All detections are included in the fitting except for NGC~4039 and NGC~3627, because the regions covered by their APEX beams are not active in star formation as indicated from their IR images. In the linear regression we assume Gaussian independent variables, and account for the errors in both \\LIR\\, and \\Lgas.  Upper limits are not adopted in the fitting. The \\Lgas-\\LIR\\, correlations extend from the nuclear regions of nearby normal galaxies to ULIRGs, covering a \\LIR\\ range of $\\sim$ 2.5 decades. We adopt publicly available IDL routines MPFIT \\citep{Markwardt2009} for the linear least-squares fit and LINMIX\\_ERR \\citep{Kelly2007} for the Bayesian regression.  The best linear least-squares fits with uncertainties are listed below, with $r$ indicating the correlation coefficient:\\\\   \\noindent \\[{\\rm log}(L_{\\rm IR})=0.95(\\pm 0.04) \\times {\\rm log}(L'_{\\rm CS   76})+4.33(\\pm 0.28);r=0.99\\] \\[{\\rm log}(L_{\\rm IR})=1.00(\\pm 0.04) \\times {\\rm log}(L'_{\\rm HCN  43})+3.67(\\pm 0.28);r=0.97\\] \\[{\\rm log}(L_{\\rm IR})=1.10(\\pm 0.05) \\times {\\rm log}(L'_{\\rm HCO^+43})+2.48(\\pm 0.38);r=0.96\\]   In the Bayesian regressions, we present the posterior distributions of possible slopes in the insets of Figure \\ref{correlations}, and obtain error estimates by fitting the distributions with Gaussian profiles. These fittings give slopes of 0.94$\\pm$0.07, 1.01$\\pm$0.07, and 1.12$\\pm$0.9 for \\CSss, \\HCNft, and \\HCOPft, respectively. These are consistent to the linear least-squares fit. Both correlations of \\LCSss-\\LIR\\ and \\LHCNft-\\LIR\\ have slopes very close to unity, while the \\LHCOPft-\\LIR\\ correlation shows a slight super-linear slope. The slopes of HCN and \\HCOP\\ are much higher than those predicted by \\citet{ncs08} and \\citet{jnm09}. Although they did not predict \\CSss, the high \\ncrit\\ of \\CSss\\ makes it far from the predicted trend too. We also fitted the correlations with our APEX data only, and obtained slopes of 0.98$\\pm$0.15, 0.98$\\pm$0.05, and 1.08$\\pm$0.05, which are very close to the results fitted with the combined data.  Alternatively, we fitted the correlations without the beam-matching correction, and obtained slopes (N) and correlation coefficients ($r$) to be: N=0.71$\\pm$0.08 and $r$=0.89 for \\CSss, N=0.77$\\pm$0.05 and $r$=0.94 for \\HCNft, and N=0.94$\\pm$0.10 and $r$=0.92 for \\HCOPft, respectively. The shallower slopes (sub-linear) and worse correlations indicate that the beam-matching correction has a significant effect in the fitting. On the other hand it should be emphasized that, with the beam-matching correction we are studying the central 14$''$-18$''$ regions in the nearby galaxies (except for the maps of \\HCNft\\ and \\HCOPft\\ in NGC 253), rather than the whole galaxies.   At \\Tkin=100K, \\HCNft, \\HCOPft, and \\CSss\\ have \\ncrit\\ of $5.6\\times 10^6$\\cmt, $1.3\\times 10^6$\\cmt, and $2.8\\times10^6$ \\cmt (see Sect \\ref{intro}), respectively. These tracers pick up the densest part of the molecular cores, and trace only a small amount of the total dense gas mass involved in star-formation. The linear slopes found in the high-$J$ HCN and CS tend to support the proposed correlations for gas tracers with \\ncrit\\ higher than $\\sim 10^4$ \\cmt\\ \\citep{Lada12}.  Comparing with the neutral HCN and CS molecules, the molecular ion \\HCOP\\ is more sensitive to the average ambient electron abundance $x(e)$, because of the protonation reaction: \\[\\rm H_3^+ + CO \\rightarrow HCO^+ + H_2\\]  The high electron abundance is likely to reduce $\\rm [HCO^{+}/H_2]$ efficiently. Considering the impacts of the CR ionization \\citep{PPP2007}, the turbulent diffusion \\citep{Xie95}, and the AGN illumination, \\HCOP\\ is likely deficit when being exposed to the extreme physical conditions prevailing at the high luminosity end, i.e., ULIRGs. Such effect could increase the slopes of the \\LHCOP-\\LIR  correlations in galaxies.   \\subsection{Comparison between Galactic Cores and Galaxies}  \\citet[][]{wes10} measured the \\LCSss-\\LIR\\ correlation in individual Galactic dense cores with high mass star formation. Their best fit yields a slope of $\\sim$ 0.8 with a large uncertainty because of the significant scatter in the data and the limited range of infrared luminosities they have at their disposal. In the bottom right of Figure \\ref{correlations}, we plot their Galactic cores and our galaxies together. Combining both samples, we find a highly linear correlation, \\[ {\\rm log}(L_{\\rm IR}) =1.00(\\pm 0.01)\\times {\\rm log}( L'_{\\rm CS\\, J=7-6})+4.03(\\pm0.04),\\] with a correlation coefficient of 0.98. This correlation is remarkably similar to that derived for HCN $J=1\\rightarrow$0 by \\citet{Wu05}, though \\CSss\\ has much higher excitation requirements. The linear correlation between \\Lgas\\ and \\LIR\\ holds over an IR luminosity range of about eight orders of magnitude.  The linearity of the (dense gas)---SFR correlation was interpreted as a ``fundamental unit'' of star formation on different physical scales, thus both SFR and dense gas mass are simply piled up by adding in more units \\citep[e.g.,][]{Wu05}. Recently it was found that once the studied gas content is denser than a density threshold, SFR does not depend on the exact value of the gas density, but depends on the total mass of the dense gas \\citep[e.g.,][]{Lada12}. In such a context the near-linearity in the $\\Sigma_{\\rm gas}-\\Sigma_{\\rm SFR}$ correlation in nearby galaxies \\citep[e.g.,][]{Bigiel2008,Schruba12} is likely caused by a constant $M_{\\rm dense}$/$M_{\\rm total\\, H_2 }$ fraction.  The gas densities probed by the three lines are beyond the highest average densities in starbursts \\citep{kt07}, and they are also one to two orders of magnitude higher than the threshold density for star-formation \\citep[\\nht $\\sim 10^4$ \\cmt; e.g., ][]{Parmentier11,Lada12}.  We find that all gases with densities $> 10^4$ \\cmt\\ are linearly correlated with SFR, which indicates that the SFR in the dense gas is not likely affected by the free-fall time scale ($t_{\\rm ff}$), because $t_{\\rm ff} \\propto \\rho^{-1/2}$. If $\\Sigma _{\\rm gas} / t_{\\rm ff}$-$\\Sigma _{\\rm SFR}$ correlations are linear for all dense gases, the shorter $t_{\\rm ff}$ for the denser gas would not keep \\Lgas-\\LIR\\ linear.  Indeed, the gas content traced by \\HCNoz\\ has 10 times longer $t_{\\rm ff}$ than that traced by \\HCNft, but they do show both linear correlations with the SFR.  The non-linear correlations found in the K-S law are likely caused by the different fractions of dense gas to total gas in different types of galaxies, rather than the free-fall timescales \\citep{Lada12}.  Actually, the K-S law with a slope of $\\sim$1.4 is likely not a fundamental physical relation, as the non-linear slope found for H$_2$+\\HI, seen in the very tenuous gas, is related to the relative amounts of dense gas and dust directly being involved in star formation and other gas being too diffuse and diluted to have anything to do with massive star formation.   \\subsection{AGN Contamination to the IR Emission}  About half of the galaxies in our sample are hosting AGNs, in particular the ULIRGs. Compared to the near- and mid IR bands, the far- and total-IR are well correlated and are thus less contaminated by the AGNs emission \\citep[e.g.,][]{U2012,jnm09}. Some AGNs, however, can still contribute significantly to the bolometric luminosity in ULIRGs \\citep[e.g.,][]{Veilleux2009}.  For the most extreme case, Mrk~231, the estimated AGN contribution is $\\sim$70\\%, and the SFR indicated from the IR emission should be $\\sim$0.5 dex lower in the \\Lgas-\\LIR\\ correlations \\citep[e.g.,][]{Veilleux2009}. For the rest galaxies, the AGN contamination should be much less ($<$10\\%) for LIRGs (\\LIR $>10^{11}$\\Lsun) and nearby Seyferts, where star-formation is expected to dominate over IR emission. Unfortunately the IR emission for most AGN-host galaxies have not been decomposed, so that we could not deduce the IR luminosity by $\\eta_{\\rm SF}$, where $\\eta_{\\rm SF}$ is the fraction of the IR luminosity due to star-formation. Overall, we do not find significant changes in the correlations if we remove AGNs from our samples. This may be due to the small size of our sample, or the fact that most AGN-hosting galaxies are not strongly contaminated by AGNs.     In conclusion, we present a survey of three dense gas tracers (\\CSss, \\HCNft, and \\HCOPft) in 20  nearby galaxies observed with the APEX 12-m telescope. Combining with data from literature and after a beam-matching correction for nearby galaxies, we find linear \\Lgas-\\LIR\\ correlations for \\CSss\\ and \\HCNft, and a slightly super-linear slope for \\LHCOPft. These results are consistent with those found in HCN\\ $J$=1$\\rightarrow$0 \\citep[i.e.,][]{gs04b,gs04a} and \\CSff\\ \\citep[i.e.,][]{Wang2011,wes10}, but contradictory to the predictions in \\citet{bns08} and \\citet{jnm09}. We also find that the linear \\LCSss-\\LIR\\ correlation can be traced $universally$ across eight orders of luminosity magnitude, down to Galactic cores. If the beam-matching corrections are not applied for nearby galaxies, however, the slopes are significantly sub-linear. Future ALMA surveys of more dense gas tracers as well as other transitions over large samples of galaxies are necessary to consolidate the above findings and ensure that, apart from the vagaries of HCO$^{+}$ chemistry, the simplest (dense gas)-SFR empirical relation is indeed true."
    },
    "1402/1402.6043_arXiv.txt": {
        "abstract": "We present the observation of a major solar eruption that is associated with fast sunspot rotation. The event includes a sigmoidal filament eruption, a coronal mass ejection, and a GOES X2.1 flare from NOAA active region 11283. The filament and some overlying arcades were partially rooted in a sunspot. The sunspot rotated at $\\sim$10$^\\circ$ per hour rate during a period of 6 hours prior to the eruption. In this period, the filament was found to rise gradually along with the sunspot rotation. Based on the HMI observation, for an area along the polarity inversion line underneath the filament, we found gradual pre-eruption decreases of both the mean strength of the photospheric horizontal field ($B_h$) and the mean inclination angle between the vector magnetic field and the local radial (or vertical) direction. These observations are consistent with the pre-eruption gradual rising of the filament-associated magnetic structure. In addition, according to the Non-Linear Force-Free-Field reconstruction of the coronal magnetic field, a pre-eruption magnetic flux rope structure is found to be in alignment with the filament, and a considerable amount of magnetic energy was transported to the corona during the period of sunspot rotation. Our study provides evidences that in this event sunspot rotation plays an important role in twisting, energizing, and destabilizing the coronal filament-flux rope system, and led to the eruption. We also propose that the pre-event evolution of $B_h$ may be used to discern the driving mechanism of eruptions. ",
        "introduction": "Solar eruptions, including solar flares, coronal mass ejections (CMEs), and filament eruptions, are spectacular energy release phenomena that occur in the solar atmosphere. They often lead to catastrophic impacts on the near-Earth space environment. They are generally believed to be a result of the rapid release of magnetic energy stored in highly-stressed/twisted magnetic structures of the corona (e.g., Forbes 2000; Low 2001). The magnetic energy is transported into the corona via slow photospheric footpoint motions, e.g., emergence, shearing, twist, etc., in a relatively long period of time, comparing to the time scale of an eruption. Among various forms of photospheric motions, sunspot rotation, first observed a century ago by Evershed (1910), has been considered to be an important process and has been studied extensively (e.g., Stenflo 1969; Barnes \\& Sturrock 1972; Ding et al. 1977, 1981; Amari et al.  1996; Tokman \\& Bellan 2002; T$\\ddot{o}$r$\\ddot{o}$k \\& Kliem 2003; Brown et al. 2003; Regnier et al. 2006; Yan et al. 2007, 2008a, 2008b, 2009; Su et al. 2010).  Previous studies confirmed the important role played by sunspot rotation in transporting energy and helicity from below the photosphere into the corona with quantitative calculations (e.g., Kazachenko et al. 2009; Vemareddy et al. 2012), and revealed some temporal and spatial association of sunspot rotation with solar flares on the basis of observational data analysis (e.g., Zhang, Li \\& Song 2007; Zhang, Liu \\& Zhang 2008; Yan et al. 2007, 2008a, 2008b, 2009; Jiang et al. 2012). There also exist a number of magnetohydrodynamic (MHD) studies examining the consequence of twisting a flux rope structure which is confined by overlying magnetic arcades (e.g., Amari et al. 1996; T$\\ddot{o}$r$\\ddot{o}$k et al. 2003). In a latest study, T$\\ddot{o}$r$\\ddot{o}$k et al. (2013) examined the role of twisting the overlying arcades in the onset of a CME using a flux rope model. These studies showed that the CME can be triggered by twisting either the core flux rope structure or the overlying coronal fields, thus established the importance of sunspot rotation in the eruption process from a theoretical perspective. On the other hand, observational studies connecting sunspot rotation with CMEs remain elusive.  Here we present a case study of the evolution of a sigmoidal filament which has roots in a rotating sunspot. The study, involving multi-wavelength imaging and vector magnetic field data from the Solar Dynamics Observatory (SDO), provides a rare case revealing the role of sunspot rotation being as not only a general energy transport process but also a direct driving process that leads to the eventual flare, CME, and filament eruption. ",
        "conclusions": "We present observations of a sunspot rotation before a major solar event consisting of a fast CME, an X2.1 flare, and a filament eruption. We suggest that this pre-eruption rotation is not only transporting energy to the corona, but also playing a dynamic and critical role in leading to the eruption. Our suggestion is based on the data analysis results, which are summarized below. Firstly, the sunspot rotation was the dominant motion in the 6 hours before the flare. The rotation rate was $\\sim$10$^\\circ$ per hour, considerably faster than some previous observations (e.g., Zhang, Liu, \\& Zhang 2008). Secondly, the filament and part of the overlying arcades were rooted in the rotating sunspot, and the filament exhibited an apparent gradual rising motion along with the sunspot rotation. This provides a possible dynamical link between the sunspot rotation and the filament dynamics as well as the resultant eruption. Thirdly, the evolutions of both the photospheric horizontal field and the magnetic field inclination angle agree with the gradual rising of the magnetic structure that supports the filament. Last, using the NLFFF method of coronal magnetic field reconstructions, we find the presence of a well-developed twisted flux rope structure associated with the filament and a considerable amount of magnetic energy increase during the sunspot rotation period. These results highlight the importance of sunspot rotation to the energy storage and the onset of the eruption.  The evolution of photospheric magnetic field is essential to both the energy build-up and the triggering of a solar eruption. Many studies have focused on rapid changes of $B_h$ induced by the flare (e.g., Wang et al. 1994, 2010; Liu et al. 2012; Wang et al. 2012). However, the detailed pre-flare evolution of this field component has not received sufficient attention. Such evolution would carry important information of the energy storage and eruption onset process. In our study, we found that $B_h$ in the area underneath the filament decreased gradually in hours before the flare. This is related to the gradual ascending of the filament-flux rope structure. Our analysis indicates that this is associated with the rapid sunspot rotation. On the contrary, studies of another active region (AR11158) revealed that $B_h$ there increased gradually in a similar time period prior to the flare (Liu et al. 2012; Wang et al. 2012). Those studies deduced that the corresponding eruptions were driven by tether-cutting reconnection (Moore et al. 2001) of two approaching magnetic loops. The pre-flare footpoint separation of these loops can explain the gradual increase of $B_h$. By comparing the different behavior of $B_h$ and corresponding understanding of the eruption mechanism, we suggest that the pre-flare variation of $B_h$ can be taken as a clue to discern the eruption mechanism: A gradual decrease of $B_h$ may be a precursor for an eruption in terms of the flux rope instability, while an increase of $B_h$ may be the precursor for tether-cutting reconnection. This needs further clarifications in future studies."
    },
    "1402/1402.7208.txt": {
        "abstract": "{Powerful radio galaxies show evidence of ongoing active galactic nuclei (AGN) feedback, mainly in the form of fast, massive outflows. But it is not clear how these outflows affect the star formation of their hosts. } { We investigate the different manifestations of AGN feedback in the evolved, powerful radio source \\tcd\\  and their impact on the molecular gas of its host galaxy, which harbors young star-forming regions and fast outflows of \\hi\\ and ionized gas. %created by the radio jet, % which could quench the star formation of the galaxy. } { We study the distribution and kinematics of the molecular gas of \\tcd\\  using high spatial resolution observations of the \\couc\\ and \\codu\\ lines, and the 3~mm and 1~continuum taken with the IRAM Plateau de Bure interferometer. We mapped the molecular gas of \\tcd\\ and compared it with the dust and star-formation images of the host.  We searched for signatures of outflow motions in the CO kinematics, and re-examined the evidence of outflowing gas in the \\hi\\ spectra. We also derived the star formation rate (SFR) and  star formation efficiency (SFE) of the host with all available SFR tracers from the literature, and compared them with the SFE of young and evolved radio galaxies and normal star-forming galaxies. } {The \\couc\\ emission line shows that the molecular gas in \\tcd\\ is distributed along a massive (\\mhd$\\sim$2.2$\\times$10$^{10}$~\\msun)  $\\sim$24$\\arcsec$(21~kpc-) diameter warped disk, that rotates around the AGN. Our data show that the dust and the star formation are clearly associated with the CO disk.  The \\codu\\ emission is located in the inner 7~kpc (diameter) region around the AGN, coincident with the inner part of the \\couc\\ disk. Both the \\couc\\ and \\codu\\ spectra reveal the presence of an absorber against the central regions of \\tcd\\ that is associated with the disk.  We do not detect any fast ($\\gtrsim$500~\\kms) outflow motions in the {  cold molecular gas}. %or warm molecular gas, the latter traced by infrared H$_2$ lines. %Comparing the \\hi\\ absorption with the \\couc\\ emission, we find that only the outer $\\sim$400~\\kms\\ of the \\hi\\ outflow component fall outside the rotating disk velocities. %, constraining the rate of outflowing neutral gas to $\\dot{M}\\simeq35$~\\msyr. The host of \\tcd\\ shows an SFE consistent with the Kennicutt-Schmidt law of normal galaxies and young radio galaxies, and it is 10-50 times higher than the SFE estimated with the 7.7~\\mum\\ PAH emission of evolved radio galaxies. Our results suggest that the apparently low SFE of evolved radio galaxies may be caused by an underestimation of the SFR and/or an overestimation of the molecular gas densities in these sources. } { The molecular gas of \\tcd, while not incompatible with a mild AGN-triggered flow, does not reach the high velocities ($\\gtrsim$500 \\kms) observed in the \\hi\\ spectrum. We find no signatures of AGN feedback in the molecular gas of \\tcd. %The molecular gas of \\tcd\\ shows no signs of ongoing AGN feedback. % The star formation and nuclear activity histories of \\tcd\\ support that the relative ages of the star forming regions and of the radio source play an important role in the impact of AGN feedback on the host galaxy. % the evolutionary scenario of AGN feedback proposed in our previous work. %The evolutionary scenario of AGN feedback proposed in our previous work does not apply to the evolved radio galaxy 3C 293. } ",
        "introduction": "%\\subsection{AGN Feedback}  %The triggering of activity in the nucleus of a galaxy can completely alter its evolution. %The presence of an active galactic nucleus (AGN) can completely alter the evolution of the host galaxy. An active galactic nucleus (AGN) releases vast amounts of energy into the interstellar medium (ISM) of its host galaxy. This energy transfer, known as AGN feedback, may increase the temperature of the ISM gas, preventing its collapse, and/or expel it in the form of outflows {  \\citep[see][ for a review]{Fabian12}}. Both effects will inhibit the star formation of the host and affect its evolution. %This transfer of energy, known as AGN feedback, is able to heat the gas in the ISM, and expel it in the form of winds (outflows). %and stop the star-formation of the host galaxy. AGN feedback can heat gas in the ISM, preventing its collapse and inhibiting star formation. It may also expel the gas of the ISM in the form of winds (outflows), which deplete the host of star-forming fuel. %Active galactic nuclei (AGN) release vast amounts of energy into the interstellar medium (ISM) of their host galaxies. This energy input can heat gas in the ISM, preventing its collapse and inhibiting star formation. It may also expel the gas of the ISM in the form of winds (outflows), which deplete the host of star-forming fuel. This transfer of energy from the AGN to the host is known as AGN feedback. This could be responsible for the correlations between black hole and host galaxy bulge mass \\citep{Magorrian98, Tremaine02, Marconi03, Haring04}. It can also explain the fast transition of early-type galaxies from the blue cloud to the red sequence \\citep{Schawinski07, Kaviraj11}. Observational evidence of the effects of AGN feedback on the hosts is frequently found in the literature \\citep{Thomas05, Murray05, Schawinski06, Muller06, Feruglio10, Crenshaw10, Fischer11, Villar11, Dasyra11, Sturm11, Maiolino12, Aalto12, Walter13, Silk13}. It is thought that AGN feedback might be responsible for several properties observed in galaxies, such as the correlations between black hole and host galaxy bulge mass \\citep{Magorrian98, Tremaine02, Marconi03, Haring04}, or the fast transition of early-type galaxies from the blue cloud to the red sequence \\citep{Schawinski07, Kaviraj11}. %Yet, observational evidence of AGN feedback is still fragmentary. Hence, it is now clear that models of galaxy evolution need to consider the effects of AGN feedback in the host galaxy \\citep{King03, Granato04, Matteo05, Croton06, Ciotti07, Menci08, King08, Merloni08, Narayanan08, Silk10}.  %Over the last decade, AGN feedback has been increasingly included in models of galaxy evolution \\citep{King03, Granato04, Matteo05, Croton06, Ciotti07, Menci08, King08, Merloni08, Narayanan08, Silk10}. {  Observational searches for signatures of AGN feedback have also grown in number} \\citep{Thomas05, Murray05, Schawinski06, Muller06, Feruglio10, Crenshaw10, Fischer11, Villar11, Dasyra11, Sturm11, Maiolino12, Aalto12}. Yet, observational evidence of AGN feedback are still fragmentary.  The AGN feedback may take place through two mechanisms: through AGN radiation (known as radiative or quasar mode of AGN feedback), or through interaction of the jet with the ISM (known as kinetic or radio mode). %%The energy transfer from the AGN to the host may take place through two mechanisms: AGN radiation (known as radiative or quasar mode of AGN feedback), or through interaction of the jet with the ISM (known as kinetic or radio mode). %the {radiative} or {quasar} mode, and the  {kinetic} or {radio} mode. In the quasar mode, the radiation from the AGN dominates the energy transfer to the ISM. In the radio mode, the momentum of the jet is transferred to the ISM. %(ARTICULO DE FARRAH!!, Review y Kaviraj tambien!) %The AGN feedback takes place through two mechanisms: the {\\it radiative} or {\\it quasar} mode, and the  {\\it kinetic} or {\\it radio} mode. In the quasar mode, the radiation from the AGN dominates the energy transfer to the ISM. In the radio mode, the momentum of the jet is transferred to the ISM. %(ARTICULO DE FARRAH!!, Review y Kaviraj tambien!) The most intense effects of the kinetic mode are seen in powerful radio galaxies. The extremely large and fast jets of powerful radio galaxies are able to inject enormous amounts of energy into their host galaxy ISM and even into the intergalactic medium \\citep[IGM; e.g.,][]{Birzan04,Fabian06,McNamara07,McNamara13,Russell13}. Most massive galaxies will harbor a radio galaxy phase during their lifetime \\citep[e.g.,][]{Best06}. Studying AGN feedback during this phase is thus crucial for understanding galaxy evolution.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} \\centering \\includegraphics[width=1.5\\columnwidth,angle=0]{c1mmc3mm_paper.eps} \\caption{Continuum maps of \\tcd\\ at 3~mm (left) and 1~mm (right). The positions of the AGN and jet components, based on the UV-FIT results (Table \\ref{uvfits}), are identified with crosses (+). %Crosses (+) mark the position of the AGN and jet peak emission according to the UV-FIT models (Table \\ref{uvfits}). The dashed box marks the size of the 1 mm map shown in the right panel. The contour levels are  --9, 9, 39 to 249 mJy~beam\\mone, in steps of 30 mJy~beam\\mone\\  for the 3~mm map; and --6, 6, 12 to 60 mJy~beam\\mone, in steps of 6 mJy~beam\\mone\\ for the 1~mm map. %$1\\sigma=2.4$ and 1.9 mJy~beam\\mone\\ for the 3 mm and 1 mm maps respectively. ($\\Delta\\,\\alpha$, $\\Delta\\,\\delta$)=(0,0) corresponds to the position of the AGN at 1 mm. The gray ellipses in the top-right corner show the PdBI beam for each map. \\label{contfits}} \\end{figure*} %1mm: lev  -3*sig_cont1 3*sig_cont1 to 30*sig_cont1 by 3*sig_cont1 %3mm: lev -3*sig_cont3 3*sig_cont3 to 90*sig_cont3 by 10*sig_cont3 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   A large number of radio galaxies show manifestations of ongoing AGN feedback at different redshift ranges, whether as gas outflows or as inhibition of star formation. Outflows in radio galaxies have been detected through \\hi\\ and ionized gas observations \\citep[e.g.,][]{Morganti03c, Rupke05, Holt06, Nesvadba06, Nesvadba08, Lehnert11}. Searches for molecular gas, which may dominate the mass/energy budget of the wind, show evidence of outflows only in one radio source  \\citep[4C~12.50;][]{Dasyra11}. %The most dramatic effects of the kinetic mode are seen in powerful radio galaxies. With jets up to several Mpc--size, they can inject enormous amounts of energy, not only into their host galaxy ISM, but also into the intergalactic medium (IGM) of the galactic group or cluster where they reside \\citep[e.g.,][]{Birzan04,Fabian06,McNamara07,McNamara12}. It is thought that a significant fraction of massive galaxies undergo a radio galaxy phase at least once in their lifetime \\citep[e.g.,][]{Best06}. Consequently, understanding how AGN feedback works during this phase is crucial for a better comprehension of galaxy evolution. %Observations show that a %A large number of radio galaxies, at different redshift ranges, host massive \\hi\\ and ionized gas outflows,  \\citep[e.g.,][]{Morganti03c, Rupke05, Holt06, Nesvadba06, Nesvadba08, Lehnert11}, but so far, only one radio source shows evidence of a molecular gas outflow  \\citep[4C~12.50;][]{Dasyra11}. %Even though molecular gas may dominate the mass/energy  budget of the wind, unambiguous evidence of a molecular outflow has  thus far only been found in one radio galaxy \\citep[4C~12.50;][]{Dasyra11}. %The inhibition of star formation, another manifestation of AGN feedback, has also been claimed to be at work in radio galaxies. Based on the molecular gas contents, and SFR estimations using the 7.7 $\\mu$m polycyclic aromatic hydrocarbons (PAH) emission, \\citet{Nesvadba10} found that radio galaxies can be $\\sim$10--50 times less efficient in forming stars than galaxies that follow the canonical Kennicutt-Schmidt \\citep[KS,][]{Schmidt59, Kennicutt98} relationship. The warm-\\hd\\ emission of these radio galaxies suggests shocks in the ISM that are induced by the expansion of the powerful radio jet. These shocks are thought to increase the turbulence in the molecular gas, which inhibits the star formation in the host galaxy. However, it is not clear how jet-induced shocks affect the cold phase of the molecular gas.  Our recent work on \\object{3C~236} \\citep{Labiano13} found that not all radio galaxies show a low star formation efficiency (SFE). \\object{3C~236} is a re-activated radio source that shows an SFE consistent with the KS-law. {  The young age of its second period of  radio-source activity \\citep[$10^5$ yr, based on the dynamics of the source,][]{Tremblay10,O'Dea01}, and the older age of the \\citet{Nesvadba10} radio source sample ($10^7-10^8$ yr) } suggested that the impact of the AGN feedback on the star formation rate probably evolves with time. A literature search showed that different authors had measured the SFR and \\mhd\\ of the young sources: \\object{4C~12.50}, \\object{4C~31.04} \\citep{Burillo07,Willett10, Dasyra12}. These two sources show an SFE consistent with the KS-law, which supports the evolutionary scenario of the feedback impact. %The different tracers used  to estimate the SFR in the young and evolved sources could also explain the discrepancy in SFE. %Another explanation for the SFE discrepancy could be due to the use of different SFR tracers. %For \\object{3C~\u00ca236} we used a large combination of tracers to evaluate its SFR, instead of just the 7.7 while the \\cite{Nesvadba10} sample was the 7.7 \\mum\\ PAH feature. % It is known that the emission of 7.7 \\mum\\ PAH may underestimate the SFR in radio galaxies, as seen in young sources \\citep[see also][]{Diamond10}. Therefore, the SFE of the evolved radio galaxies could be higher than measured, even consistent with the KS-law, and the proposed evolutionary scenario would not apply. %{ ESCRIBIR BIEN: It could also be that the SFR tracers used (7.7um) are not addequate for radio galaxies/underestimate the SFR in radio galaxies. Nuestro trabajo en 3C236, una recently restarted RG, encontro que la SFR era consistente con la KS-law. Podria ser que los SFR tracer de la N10 no sean buenos o que los efectos del feedback necesiten tiempo para propagarse por el ISM, ya que tb vimos que las de Willet (4C...) eran consistentes con las normales.} Observations of cold molecular gas in the ISM of different populations of young and evolved radio galaxies and accurate measurements of their SFR are the key to understanding how AGN feedback affects the properties of their hosts.  %\\citet{Nesvadba10} used the 7.7 $\\mu$m polycyclic aromatic hydrocarbons (PAH) emission to estimate the star-formation rate (SFR) in a sample of radio galaxies. Based on their estimated molecular gas contents,  \\citet{Nesvadba10} found that radio galaxies in their sample are $\\sim$10--50 times less efficient in forming stars compared to normal galaxies that follow the canonical Kennicutt-Schmidt relationship \\citep{Schmidt59, Kennicutt98}. %The warm-\\hd\\ emission of these radio galaxies suggests the presence of shocks in the warm molecular gas traced by a set of infrared H$_2$ lines. These shocks, identified in the warm-H$_2$ phase,  would increase the turbulence in the molecular ISM as a whole, inhibiting the star formation in the host. Nevertheless, it is still an open question whether shocks are affecting the bulk of the molecular ISM or if they only concern the  warm molecular gas phase. CO observations, which are well adapted to trace most of the molecular ISM, are the key to understanding how AGN feedback affects the kinematics and the star-formation properties of molecular gas in radio galaxies.  %Observations of \\hi\\ and ionized gas in radio galaxies have found massive outflows in a significant number of radio galaxies at different redshift ranges \\citep[e.g.,][]{Morganti03c, Rupke05, Holt06, Nesvadba06, Nesvadba08, Lehnert11}. Even though molecular gas may dominate the mass/energy  budget of the wind, unambiguous evidence of a molecular outflow has  thus far only been found in one radio galaxy \\citep[4C~12.50;][]{Dasyra11}. Another manifestation of AGN feedback, the inhibition of star formation, has also been claimed to be at work in radio galaxies. \\citet{Nesvadba10} used the 7.7 $\\mu$m polycyclic aromatic hydrocarbons (PAH) emission to estimate the star-formation rate (SFR) in a sample of radio galaxies. Based on their estimated molecular gas contents,  \\citet{Nesvadba10} found that radio galaxies in their sample are $\\sim$10--50 times less efficient in forming stars compared to normal galaxies that follow the canonical Kennicutt-Schmidt relationship \\citep{Schmidt59, Kennicutt98}. The warm-\\hd\\ emission of these radio galaxies suggests the presence of shocks in the warm molecular gas traced by a set of IR H$_2$ lines. These shocks, identified in the warm-H$_2$ phase,  would increase the turbulence in the molecular ISM as a whole, inhibiting the star formation in the host. However,  it is still an open question whether shocks are affecting the bulk of the molecular ISM or if they only concern the  warm molecular gas phase. CO observations, which are well adapted to trace most of the molecular ISM, are the key to understanding the {complexity} of how AGN feedback affects the kinematics and the star-formation properties of molecular gas in radio galaxies.      \\subsection{\\tcd}  %ESTRUCTURA DE RADIO The radio source \\tcd\\ is a large ($\\sim$200 kpc), powerful ($P_{\\rm 5GHz}$=9$\\times$$10^{24}$ W~Hz\\mone), %double-double \\citep{Schoenmakers00}, FR II\\footnote{\\tcd\\ also shows properties of an FR I radio galaxy. In the literature, it is sometimes classified as an FR I or intermediate FR I/II radio galaxy \\citep[e.g.,][]{Leipski09, Massaro10, Tadhunter11}.}  \\citep{Fanaroff74} radio galaxy with PA=135\\grados. The inner 2$\\arcsec$ of the radio source show a compact core and two jets aligned east-west (PA=93\\grados), with the eastern jet approaching. %This inner structure is disconnected from the large-scale structure \\citep[e.g.,][]{Bridle81, Breugel84, Akujor96, Beswick02, Beswick04, Giovannini05}. The double-double structure suggests an episodic activity of the nuclear activity \\citep[e.g.,][]{Saikia09}. Estimations based on spectral aging show that the large radio source was triggered $\\lesssim$20 Myr ago and stopped $\\sim$0.35 Myr ago, while the compact source was triggered $\\lesssim$0.18 Myr ago \\citep{Akujor96, Joshi11}. %; the second one been triggered by the merger \\citep{Tadhunter11}. %Assuming a constant hotspot advance speed of 0.1$c$, \\citet{Akujor96} estimate ages 3$\\times$10$^4$ yr and 2.7 Myr for the inner and outer radio sources respectively. The time-lapse between the two activity epochs probably lasted only  $\\lesssim$0.1 Myr \\citep{Joshi11}. This short pause in activity implies that the radio jets of \\tcd\\ may have been almost continuously interacting with the ISM of the host for the last $\\sim$20 Myr.  %HOST The host galaxy of \\tcd\\ \\citep[\\object{UGC~8782} or \\object{VV5-33-12},  $z$=0.0450,][]{Vaucouleurs91} has been classified both as a spiral galaxy \\citep[e.g.,][]{Sandage66, Colla75, Burbidge79} and as an elliptical galaxy \\citep[e.g.,][]{Veron01, Tremblay07}, probably because of its extremely complex morphology, with a dust disk, compact knots, and large dust lanes\\footnote{The total mass of dust in \\object{UGC~8782} is $\\sim$$10^7$ \\msun\\ \\citep{Koff00, Papadopoulos10}.}. The morphology and the vast amounts of gas and dust found in the host of \\tcd\\ are consistent with a gas-rich merger of an elliptical galaxy with a spiral galaxy \\citep{Martel99, Koff00, Capetti00, Floyd06}. \\object{UGC~8782} shows an optical emission tail that extends beyond a possible companion $\\sim$37$\\arcsec$ to the southwest \\citep{Heckman86, Smith89}. It is not clear whether this companion galaxy is interacting with \\object{UGC~8782} or was formed in the dust tail \\citep{Floyd06}.  %UV and SF Hubble Space Telescope (HST) near-UV (NUV) photometry of \\object{UGC~8782} revealed several regions of bright emission along the edges of the dust lanes, consistent with recent star formation  \\citep{Allen08, Baldi08}. Optical spectroscopy shows that the total stellar mass of \\tcd\\ is 2.8$\\times$10$^{11}$ \\msun, consisting of a young (0.1-2.5 Gyr) stellar population, which represents 57\\% of the total stellar mass, and an old stellar population \\citep[$\\gtrsim$10 Gyr,][]{Tadhunter05}. The age of the young stellar population is consistent with a star formation episode after the merger \\citep{Tadhunter05, Tadhunter11}. The {\\it Starlight} \\citep{Cid05} fit of the SDSS spectrum Ca H+K lines showed that the stellar population of \\tcd\\ can be described by three components with ages 10 Gyr, 0.1--1 Gyr, and 10 Myr \\citep{Nesvadba10}.  %The merger would also be responsible for the triggering of the act period%The young stellar population represents 57\\% of the total stellar mass \\citep{Tadhunter11} of  \\object{UGC~8782}.   The \\hi\\ spectrum of \\tcd\\ shows a broad ($\\sim$1200 \\kms)  absorption component associated with the western jet of the compact radio source (0.5 kpc west of the nucleus). This absorption is created in an outflow of neutral gas  produced by the interaction of the radio jet with the ISM \\citep{Morganti03, Morganti05, Mahony13}. Optical, long-slit spectra show an outflowing component in the ionized gas emission lines, with a width and center consistent with those of the \\hi\\ outflow. This ionized gas outflow, however, is associated with the east component of the compact radio source (1 kpc east of the nucleus). Weaker hints of outflowing ionized gas might also be present in the center and western components of the compact radio source \\citep{Emonts05}. %If the systemic velocity of \\tcd\\ were estimated based on the \\hi\\ absorption spectrum, the outflow component would appear blueshifted. If it is calculated using the emission lines, the outflow component of the spectrum is centered at the systemic velocity.  %Previous searches of molecular gas in \\tcd\\ were plagued by limited sensitivity and they were undertaken with single-dish telescopes.  Based on CO observations made with the NRAO-12m telescope and the OVRO Millimeter Array, \\citet{Evans99b} reported the detection of the \\couc\\  line, both in emission and absorption, in \\tcd\\  \\citep[see also][]{Evans05}. According to their data, which have a spatial resolution of  $\\sim$3.5$\\arcsec$, the \\couc\\ emission extends across a 7$\\arcsec$ diameter disk, with a total mass of cold \\hd\\ of  \\mhd$\\sim$1.5$\\times$10$^{10}$ \\msun, and is perpendicular to the large-scale radio source. %Their velocity coverage ($\\sim$1200 \\kms) and sensitivity (0.012 Jy) did not allow a detection of a cold molecular gas outflow. %, which is higher, by almost an order of magnitude, than the cold-\\hd\\ masses usually found in re-started radio galaxies \\citep[e.g.,][]{Saripalli07}. Later observations, using the JCMT, detected high-excitation CO emission lines in \\tcd\\ \\citep[\\coct, \\cosc,][]{Papadopoulos08, Papadopoulos10}. %, which could have been enhanced by the outflow. %, and showed no presence of a molecular gas outflow either. %  The outflows could be enhancing the emission of the high excitation CO lines detected \\citep{Papadopoulos08}. The PdBI was also used to probe the molecular gas using the \\codu\\  line, and search for evidence of outflow \\citep{Burillo09}. None of the observations above had the required velocity coverage and sensitivity to detect the outflow seen in \\hi\\ and ionized gas, however.  Even though \\tcd\\ has large contents of molecular gas and young stars, \\citet{Nesvadba10} measured an SFE 20 times lower than expected for normal star-forming galaxies, based on its 7.7 \\mum\\ PAH emission and the \\hd\\ measurements of \\citet{Evans99b}. They attributed this low SFE to the impact of AGN feedback on the ISM through the outflows observed in \\hi\\ and ionized gas. On the other hand, the FIR luminosity of \\tcd\\  \\citep[$3.3\\times10^{10}$~\\lsun,][]{Floyd06} is consistent with FIR luminosities of normal, KS-law, star-forming galaxies with similar molecular gas contents, suggesting that the SFR of \\tcd\\ calculated with the 7.7 \\mum\\ PAH emission may be underestimated. %, yields a SFE consistent with normal star-forming galaxies with similar molecular gas contents. %{ Solving this discrepancy in SFR is key .... } %It is therefore essential to obtain an accurate measurement of the SFR of \\tcd\\ to understand the impact of AGN feedback on the star formation of its host.   %Based on its FIR luminosity \\citep[3.3\\times10^{10}$\\lsun,][]{Floyd06} and molecular gas mass, \\tcd\\ seems to be forming stars at a normal, KS-law, efficiency. On the other hand, %indicative of star formation, is 3.3\\times10^{10}$\\lsun\\ \\citep{Floyd06}, consistent with FIR luminosities of  normal star-forming galaxies with similar molecular gas contents. On the other hand, %Even though the FIR luminosity of \\tcd\\ is $\\sim$3 times lower than Luminous Infrared Galaxies (LIRGs) with similar molecular gas contents, it is consistent with normal star-forming galaxies. On the other hand, %Even though \\tcd\\ has large molecular gas and young star contents, its far-infrared (FIR) luminosity ($L_{\\mathrm{FIR}}=3.3\\times10^{10}$\\lsun), indicative of star formation, is low compared to Luminous Infrared Galaxies (LIRGs, L$_{\\mathrm{FIR}}\\gtrsim10^{11}$\\lsun) with similar molecular gas masses \\citep{Floyd06}. %\\citet{Nesvadba10} measured a star-formation efficiency 20 times lower than expected for normal star-forming galaxies, based on the 7.7  \\mum\\ PAH emission of \\tcd. They claim that this low efficiency by the fast outflows..... % REESCRIBIR LO SIGUIENTE Y SEPARAR EL OUTFLOW! This low star-formation efficiency (SFE) could be caused by the fast ($\\gtrsim$1000 km~s\\mone) outflows of \\hi\\ and ionized gas detected by \\citet{Morganti03} and \\citet{Emonts05} in \\tcd. The outflows, produced by the interaction between the radio jets of the inner radio source and the ISM of the host galaxy, could also be enhancing the emission of the high excitation CO lines detected \\citep{Papadopoulos08}.  %The lack of a stellar population younger than the compact radio source suggests that the merger could have happened recently and the gas had not enough time to settle and form stars. %The low star-formation rate (SFR) of \\tcd\\ suggests that the merger was recent, and the SF has not had time to get underway properly, or that something (As AGN feedback) is suppressing the SFR. REESCRIBIR UN POCO. %Even though the young stellar population accounts for more than half of the stellar mass, the star-formation rate (SFR) of \\tcd\\ could be smaller than expected for the large amounts of molecular gas found in its ISM.   The large molecular gas content of the host, the ages of the old radio source and young stellar population, the re-ignition of the nuclear activity, and  gas outflows caused by jet-ISM interactions make \\tcd\\ an ideal candidate for studying the impact of AGN feedback in radio galaxies. %This is the third of a series of papers where we use interferometric observations of the molecular line emission to study the fueling and the feedback of activity in a sample of nearby radio galaxies \\citep[observations of \\object{4C~31.04} and \\object{3C~236} were presented by][]{Burillo07, Labiano13}. We carried out 1~mm, 3~mm continuum and \\couc, \\codu\\ line observations of \\tcd\\ with the Plateau de Bure Interferometer (PdBI)  to investigate the impact of AGN feedback of activity in its host. Our sensitivity and resolution at 3 mm are, respectively, five  and two times  better than achieved before \\citep[e.g.,][]{Evans99b}, %(160 times at 1 mm) and we covered a velocity range broad enough ($>$4000 \\kms) to detect the molecular gas counterpart of the \\hi\\ and ionized gas outflows. %To this purpose, we have carried out 1~mm, 3~mm continuum and \\couc, \\codu\\ line high-resolution and high-sensitivity observations of \\tcd\\ with the Plateau de Bure Interferometer (PdBI). %We reached a resolution 10 times higher and a sensitivity 2 times better than previous observations, covering 6000 \\kms\\ in velocity. With our data, the outflow of molecular gas, if present, will be observed to its full extent.W In this paper, we present the results of these observations where we study the continuum emission of \\tcd, and analyze the morphology and kinematics of the molecular gas. We determine the systemic velocity of \\tcd, search for outflow signatures in the {  cold molecular gas}, and compare our results with the \\hi\\ observations of the outflow. We then reassess the SFR of \\tcd\\ based on several SFR tracers, taking care to examine possible effects of the AGN. Finally, we compare the SFE of \\tcd\\ with a sample of powerful radio galaxies and normal star-forming galaxies and examine the possible impact of AGN feedback on star formation. %We mapped and studied the continuum emission of  \\tcd, and analyzed the distribution and kinematics of molecular gas based on the \\couc\\ and \\codu\\ emission lines. We accurately determined systemic velocity of \\tcd\\ and searched for outflow signatures in the cold and warm molecular gas, and compare our results with the \\hi\\ observations of the outflow. We derived the SFR of \\tcd\\  based on all the SFR tracers available from the literature and archival data to minimize the uncertainties introduced by the AGN emission or the destruction of tracers (such as the 7.7 \\mum\\ PAH) in an AGN environment. Based on the mass of molecular gas derived from our observations, and on the SFR estimations, we compared the SFE of \\tcd\\ with a sample of young and evolved powerful radio galaxies, and normal star-forming galaxies. The comparison allowed us to study the impact of AGN feedback on the star formation of the hosts of radio galaxies, and to check the applicability of the evolutionary scenario proposed in \\citet{Labiano13} as an explanation for the apparently low SFE of evolved radio galaxies. %Finally, we discuss a scenario consistent with the properties and the history of \\tcd\\ and its host.  %============\\\\  %Las cosas de de Notas.txt!!\\\\ %Poner en antedcedentes con  todo el taco de articulos que me lei al empezar.\\\\ %Check Proposal.\\\\ %PROPIEDADES DEL HOST. Eliptica? Masa de estrellas, masa total.\\\\ %HISTORIA DE LA FUENTE. Reactivacion? Mergers? Mas cosas que destacar?\\\\ %OBSERVACIONES PREVIAS DE \\tcd. Nuestras. De otros grupos?\\\\ %ANTECEDENTES DE FEEDBACK EN LA  \\tcd? Outflow de \\citet{Morganti05}, Emonts?\\\\ %PROPIEDADES UV %PROPIEDADES IR %PROPIEDADES HI %PROPIEDADES CO %Por que es interesante en esta fuente.\\\\ %This is the third of a series of papers?\\\\ ",
        "conclusions": "%The powerful radio galaxy \\tcd\\ was known to harbor large molecular gas contents, active star-forming regions and \\hi\\ and ionized gas  outflows. %\\tcd\\ is a powerful radio galaxy known to harbor a re-ignited AGN, active star-forming regions, large molecular gas contents, and a fast \\hi\\ and ionized gas outflow.  We carried out 1~mm, 3~mm continuum and \\couc, \\codu\\ line PdBI observations of the powerful radio galaxy \\tcd. The host of \\tcd\\ was known to harbor active star-forming regions and outflows on \\hi\\ and ionized gas. The  high sensitivity and resolution of our data in combination with HST and Spitzer images and \\hi\\ spectra allowed us to study the impact of AGN feedback on the host of this evolved radio galaxy in detail. We investigated the kinematics and distribution of the molecular gas to search for evidence of outflow motions. We also accurately determined the SFE of the host galaxy and compared it with the efficiencies of normal star-forming galaxies and young and old radio sources. The main results of our work are summarized below:   %with unprecedented sensitivity and resolution to investigate the impact of AGN feedback of activity in its host galaxy. \\tcd\\ is a powerful, evolved radio galaxy with active star forming regions and outflows on \\hi\\ and ionized gas. %, the age of the old radio source, which has been active until recently, the age of the young, massive, stellar population, the re-ignition of the nuclear activity, and the detection of fast outflows due jet-ISM interactions, make \\tcd\\ an ideal candidate to study the impact of AGN feedback in radio galaxies. %This is the third of a series of papers where we use interferometric observations of the molecular line emission to study the fueling and the feedback of activity in a sample of nearby radio galaxies \\citep[observations of \\object{4C~31.04} and \\object{3C~236} were presented by][]{Burillo07, Labiano13}. %We carried out 1~mm, 3~mm continuum and \\couc, \\codu\\ line observations of \\tcd\\ with the Plateau de Bure Interferometer (PdBI)  to investigate the impact of AGN feedback of activity in its host. Our resolution at 3 mm is 9 times better than achieved before (160 times at 1 mm), and we covered a velocity range large enough ($>$4000 \\kms) to detect the molecular gas counterpart of the \\hi\\ and ionized gas outflows. %To this purpose, we have carried out 1~mm, 3~mm continuum and \\couc, \\codu\\ line high-resolution and high-sensitivity observations of \\tcd\\ with the Plateau de Bure Interferometer (PdBI). %We reached a resolution 10 times higher and a sensitivity 2 times better than previous observations, covering 6000 \\kms\\ in velocity. With our data, the outflow of molecular gas, if present, will be observed to its full extent.W %We mapped and studied the continuum emission of  \\tcd, and analyzed the distribution and kinematics of molecular gas based on the \\couc\\ and \\codu\\ emission lines. We accurately determined systemic velocity of \\tcd\\ and searched for outflow signatures in the cold and warm molecular gas, and compare our results with the \\hi\\ observations of the outflow. We derived the SFR of \\tcd\\  based on all the SFR tracers available from the literature and archival data, to minimize the uncertainties introduced by the AGN emission or the destruction of tracers (such as the 7.7 \\mum\\ PAH) in an AGN environment. Based on the mass of molecular gas derived from our observations, and on the SFR estimations, we compared the SFE of \\tcd\\ against a sample of young and evolved powerful radio galaxies, and normal star-forming galaxies. The comparison allowed us to study the impact of AGN feedback on the star formation of the hosts of radio galaxies, and to check the validity of the evolutionary scenario proposed in \\citet{Labiano13}. %Finally, we discuss a scenario consistent with the properties and the history of \\tcd\\ and its host.  %Using \\couc\\ and \\codu\\  observations of the IRAM PdBI, we have studied the distribution and kinematics of the molecular gas in the radio galaxy \\tcd. This galaxy was known to harbor a re-ignited AGN, recent star-formation, large molecular gas contents, and a fast \\hi\\ and ionized gas outflow. The new observations presented in this paper were analyzed to search for the footprints of AGN  feedback on the molecular ISM of  \\tcd. In particular, we looked for evidence of outflow motions in the molecular gas kinematics. We also derived the SFE of molecular gas in \\tcd\\ and compared this value with the SFE obtained in different populations of radio galaxies, including $young$ and $evolved$ radio sources. We investigated if the star-formation properties of \\tcd\\ deviate from the KS-law followed by normal star-forming galaxies.  \\begin{itemize}  \\item  { The 1~mm and 3~mm maps separate the continuum emission from the AGN and the eastern jet of \\tcd. The emission of the AGN and jet at 1~mm and 3~mm are consistent with synchrotron radiation. %While the 3~mm emission is consistent with synchrotron emission, the 1~mm continuum shows a combination of synchrotron radiation from the AGN and the jet plus dust or free-free emission from the AGN coordinates, possibly originated in the dust lane crossing in front of the core. }  \\item  {  The cold molecular gas of \\tcd, traced by the \\couc\\ emission, forms a disk with a diameter of 21~kpc, with a mass \\mhd=2.2$\\times$10$^{10}$ \\msun. The dust mapped by the HST $R$--$H$ image and the star formation seen in the HST-NUV and Spitzer-8 \\mum\\ images are clearly associated with the molecular gas disk. % The \\codu\\ emission of \\tcd\\ is originated in a disk of diameter 7~kpc around the AGN. The morphology of the CO emission is consistent with a warped, corrugated disk, with the outer edges ($R\\gtrsim8$ kpc) seen almost edge-on, and the central ($R\\lesssim4$ kpc) region tilted toward the plane of the sky, with the southern side facing us. }  \\item  {  The kinematics of the CO disk is consistent with rotation around the core, with no evidence of fast ($\\gtrsim$500 \\kms) outflowing molecular gas. %The warm phase of the molecular gas does not show traces of outflow either. The detailed analysis of the CO kinematics showed that the inclination and the position angle of the disk are consistent with a corrugated, warped morphology. Furthermore, the Fourier decomposition of the velocity field {is consistent with} the presence of a warp in the CO disk of \\tcd, as suggested by the CO maps. %CO kinematics showed that the inclination angle is roughly constant ($i$$\\simeq$75\\grados) for radii below 3~kpc, and increases from 20\\grados\\ (at radius 3.2 kpc) to 65\\grados\\ at the edge of the disk. The position angle is roughly constant up to radii 6.4 kpc and shifts to 55\\grados\\ beyond that radius. The Fourier decomposition of the velocity field Using the kinematics of the \\couc\\ emission line, we determined the systemic velocity of \\tcd\\ to be \\vsco=13\\,434$\\pm$8~\\kms.  %8 kms es la mitad del ancho del canal de CO10 }  \\item  The CO spectra of \\tcd\\ reveal several CO absorption features toward the core and jet coordinates. These absorptions are consistent with those observed in the \\hi\\ spectra, which originate in the disk. % They seem to be originated in the dust lanes  that criss-cross in front of \\tcd. The comparison of the CO and \\hi\\ spectra also showed that the velocities of the \\hi\\ outflow that cannot be explained by rotation of the disk range from $\\sim$300 \\kms\\ to $\\sim$900 \\kms\\ with respect to the systemic velocity. % is $\\lesssim$400 \\kms, with an outflow rate $\\dot{M}\\simeq35$ \\msyr. %  This lowers the $\\dot{M}$ derived from the \\hi\\ spectra.  \\item  We studied the different SFR tracers available for \\tcd, from NUV to IR, to obtain an accurate estimate of its SFR. The most reliable tracers are the [\\ion{Ne}{ii}] and 11.3 \\mum\\ PAH emission, as well as the 24 and 70 \\mum\\ continuum, which, for \\tcd, have negligible AGN contributions. The average SFR value for \\tcd\\ is 4.0$\\pm$1.5 \\msyr.  \\item The SFE of \\tcd\\ ($SFE$=0.18~Gyr\\mone) is consistent with the KS-law of normal galaxies and young radio galaxies. It is 10-50 times higher than the efficiencies measured in other evolved radio galaxies. Therefore, the evolutionary scenario of AGN feedback proposed in our previous work \\citep{Labiano13}  does not apply to the evolved radio galaxy \\tcd. The higher SFE of \\tcd\\ is a global property of the molecular gas disk.   %\\item %The SFE of the molecular disk is consistent with the SFE of young radio sources. %Therefore, the evolutionary scenario of AGN feedback proposed in our previous work \\citep{Labiano13} is also consistent with evolved radio sources where the nuclear activity was triggered long after the starburst. %Therefore, the evolutionary scenario of AGN feedback proposed in our previous work \\citep{Labiano13}  does not apply to the evolved radio galaxy \\tcd.  The discrepancies between the SFE of radio galaxies and normal star-forming galaxies may be due to an underestimation of the SFR and/or an overestimation of the \\hd\\ densities.%   \\item The origin of the discrepancy between the SFE of powerful radio galaxies and normal, KS-law, star-forming galaxies, might be due to the use of the 7.7\\mum\\ PAH emission, which probably underestimates the SFR in powerful AGN environments, and/or an underestimation of the \\hd\\ surface densities in radio galaxies.  \\end{itemize}"
    },
    "1402/1402.2582_arXiv.txt": {
        "abstract": "The preparation of a space-mission that carries out any kind of imaging to detect high-precision low-amplitude variability of its targets requires a robust model for the expected performance of its instruments. This model cannot be derived from simple addition of noise properties due to the complex interaction between the various noise sources. While it is not feasible to build and test a prototype of the imaging device on-ground, realistic numerical simulations in the form of an end-to-end simulator can be used to model the noise propagation in the observations. These simulations not only allow studying the performance of the instrument, its noise source response and its data quality, but also the instrument design verification for different types of configurations, the observing strategy and the scientific feasibility of an observing proposal. In this way, a complete description and assessment of the objectives to expect from the mission can be derived. We present a high-precision simulation software package, designed to simulate photometric time-series of CCD images by including realistic models of the CCD and its electronics, the telescope optics, the stellar field, the jitter movements of the spacecraft, and all important natural noise sources. This formalism has been implemented in a software tool, dubbed ASTROID Simulator. ",
        "introduction": "The ASTROID Simulator is a software package for realistic modeling of high-precision high-cadence space-based imaging observations of a selected stellar field. It permits to estimate the impact of instrumental and natural noise sources to predict the quality of data and the performance of the instrument. This tool has been designed to perform simulations for a large range of different set-ups thanks to its flexibility and an extended set of input parameters; the ASTROID Simulator is built to be applicable to space missions with a wide range of detector characteristics. The simulator is currently considering charge coupled device (hereafter CCD) simulations but has been redeveloped in a modular architecture that easily allows modifications and implementation of new functionalities, so it can be used for different kinds of future missions with straightforward modifications.  In order to accomplish the multi-mission task, the simulator has been constructed based on two main ideas: employing an architecture based on modularity principles and mimicking a standard science imaging pipeline. The modularity permits to treat any of the steps in the processing independently, and add or modify the implemented functionalities. Any inexperienced user can easily improve the comprehension of the process due to the standard pipeline architecture. The standard pipeline architecture also makes the access to the source code lighter. For users who want to adapt this simulator to particular space missions, it is easy to identify whatever step in the process is different or has any different feature than those of the present standard regular processing.   \\begin{figure*}[!ht] \\centering \\includegraphics[width=\\textwidth, clip,angle=0]{Figures/1color.eps} \\caption{Architectural design components relationship overview.} \\label{fig:architecture} \\end{figure*}      A photometry module has been implemented to analyze the simulations. This module analyzes every image generated in the CCD processing module to obtain photometric fluxes and compares the input star catalogue sources with the sources obtained in the produced images. This is described in detail in Sect. 4. The ASTROID Simulator is based on the original PLATO Simulator \\citep{Zima2010} which was developed for the assessment study of the PLATO M3 mission candidate within ESA's Cosmic Vision 2015-2025 programme. The ASTROID Simulator has been redesigned and implemented as a multi-mission imaging simulator aiming the ease of including new functionalities and detectors. \\cite{Catala2008}, \\cite{Claudi2010}, \\cite{Rauer2010} provide detailed descriptions of the PLATO mission. The PLATO Simulator, in turn, is based on pre-existing codes that were developed for the ESA Eddington mission candidate (\\citealt{Arentoft2004}; \\citealt{DeRidder2006}). ",
        "conclusions": ""
    },
    "1402/1402.2311_arXiv.txt": {
        "abstract": "We analyse three cosmological scenarios with interaction in the dark sector, which are particular cases of a general expression for the energy flux from vacuum to matter. In the first case the interaction leads to a transition from an unstable de Sitter phase to a radiation dominated universe, avoiding in this way the initial singularity. In the second case the interaction gives rise to a slow-roll power-law inflation. Finally, the third scenario is a concordance model for the late-time universe, with the vacuum term decaying into cold dark matter. We identify the physics behind these forms of interaction and show that they can be described as particular types of the modified Chaplygin gas. ",
        "introduction": "The interaction between the two components of the dark sector of the Universe is not only allowed by Einstein's equations but also leads to interesting scenarios for both the primordial and late-time phases of the expansion \\cite{chimento1,zimdahl,sergio,elcio}. In this letter we review three of these scenarios and show that they are particular cases of a general expression for the energy flux from vacuum to the matter sector. In the first of them, the vacuum-matter interaction leads to a transition from an unstable de Sitter phase to a radiation dominated universe, avoiding in this way the initial singularity. In the second case, the interaction gives rise to an inflationary expansion without inflaton, with desired features of standard inflationary models. Finally, the third scenario is a concordance model for the late-time universe, with the vacuum term decaying into cold dark matter. In the three cases we identify the physics behind these particular forms of interaction. We also show that they can be interpreted as three particular types of the modified Chaplygin gas. ",
        "conclusions": "The three cases above can be written in a unified form. Indeed, by using the pairs (\\ref{Q1})-(\\ref{H4}), (\\ref{Q2})-(\\ref{inflation}) and (\\ref{Q3})-(\\ref{linear}), we can write the energy flux as \\begin{equation} \\label{Q} Q = - \\alpha \\,\\frac{\\rho_x \\rho'}{\\rho}, \\end{equation} where $\\alpha = -2, -1, -1/2$ in the first, second and latter cases, respectively. On the other hand, from (\\ref{Friedmann}) and (\\ref{conservation2}) it is easy to show that \\begin{equation} \\rho_x = \\rho + \\frac{\\rho'}{\\gamma_m}. \\end{equation} Substituting this result into (\\ref{Q}), we obtain \\begin{equation} Q = - \\alpha \\left( \\rho' + \\frac{\\rho'^2}{\\gamma_m \\rho} \\right). \\end{equation} This last equation has the general form of what has been called a homogeneous non-linear interaction in the dark sector \\cite{chimento2}.  It is also noteworthy that in the three cases the dark sector can be described in a unified way as a modified Chaplygin gas. Indeed, equations (\\ref{H4}), (\\ref{inflation}) and (\\ref{linear}) can be put in the general form \\begin{equation} \\rho_x = A \\rho^{-\\alpha}, \\end{equation} with $\\alpha$ defined as above. Now, by using \\begin{eqnarray} \\rho &=& \\rho_x + \\rho_m,\\\\ p &=& -\\rho_x + (\\gamma_m -1) \\rho_m, \\end{eqnarray} we can show that \\begin{equation} p = (\\gamma_m -1) \\rho - \\frac{B}{\\rho^{\\alpha}}, \\end{equation} with $B = \\gamma_m A$. In the late-time case of non-relativistic dark matter ($\\gamma_m = 1$), the dark sector is  described by a generalised Chaplygin gas with negative $\\alpha$ \\cite{fabris}, also named extended Chaplygin gas \\cite{chimento3}."
    },
    "1402/1402.3121_arXiv.txt": {
        "abstract": "We have solved the Einstein-Maxwell equations for a class of metrics with constant spatial curvature by considering only a primordial magnetic field as source. We assume a slight modification of the Tolman averaging relations so that the energy-momentum tensor of this field possesses an anisotropic pressure component. This inhomogeneous magnetic universe is isotropic and its time evolution is guided by the usual Friedmann equations. In the case of a flat universe, the space-time metric is free of singularities (except the well-known initial singularity at t = 0). It is shown that the anisotropic pressure of our model has a straightforward relation to the Weyl tensor. We then analyze the effect of this new ingredient on the motion of test particles and on the geodesic deviation of the cosmic fluid. ",
        "introduction": "The presence of magnetic fields in the early universe is fundamental to explain the evolution of the large scale structures of the universe. The current understanding of structure formation indicates that, beyond dark matter, magnetic fields are necessary to form galaxies and stars in accordance to what we observe (for a review see \\cite{durrer} and references therein). In particular, magnetic fields are currently studied in the context of small perturbations of the standard cosmological model \\cite{barrow,pogosian} with possible deviations caused by its presence \\cite{caprini,shaw}. % On the other hand, aiming to produce a more realistic description of the universe, some authors claim that the main problem of the standard cosmological model (SCM) lies on the huge simplification of the Friedmann-Lema\\^itre-Robertson-Walker (FLRW) geometry adopted as a background. Hence, in a variety of ways, they suggest modifications of spacetime symmetries leading to inhomogeneous cosmological models \\cite{lemaitre,tolman,bondi,szekeres,mcvittie,wiltshire} (for a review see \\cite{ltb}). In both cases (SCM and inhomogeneous models), magnetic fields are widely studied and should have interesting consequences, for example the avoidance of initial singularity \\cite{delorenci02,aline}.  We consider in this work the evolution of a cosmological model in the presence of a primordial magnetic field. In addition to that, we slightly modify the Tolman relations for the averages of the electromagnetic field, and that renders our model inhomogeneous. The effective energy-momentum distribution of the magnetic field in this case possesses an anisotropic pressure term $\\pi_{\\mu\\nu}$ that provides non-trivial properties for the metric.  As mentioned in Ref.\\ \\cite{ltb}, $\\pi_{\\mu\\nu}$ has been indeed considered in the establishment of alternative models, but it is ruled out in standard cosmology due to the violation of the FLRW symmetries and its association restricted to the shear tensor as dictated by the fluid mechanics \\cite{landau}. However, in this work we show that the Einstein-Maxwell equations allow a set of solutions which is driven by the usual Friedmann equations and has a 3-space of constant spatial curvature, given by the Schwarzschild-de Sitter lattice \\cite{rindler}. The source of these equations may be represented by a radiation fluid ($p=\\rho/3$) with an extra anisotropic pressure component. In contrast to most inhomogeneous cosmological models, ours is part of the shear-free solutions of Einstein equations \\cite{ellis}. Therefore, $\\pi^{\\mu}{}_{\\nu}$ should be linked to another symmetric traceless quantity which in this case it is shown to be the electric part of the Weyl tensor.  This work is organized as follows. In Sec.\\ \\ref{II} we introduce the model with the stochastic magnetic field as source for the Einstein equations and see the consequences of the presence of an anisotropic pressure component due to this field on the evolution equations for the scale factor. We then analyze qualitatively the case of a flat space and the relation between $\\pi^{\\mu}{}_{\\nu}$ and the Weyl tensor. In Sec.\\ \\ref{III} we analyze the kinematical properties of the model through the study of geodesic motion and geodesic deviation. We finally present our conclusions in Sec.\\ \\ref{IV}. ",
        "conclusions": "\\label{IV}  From an exact solution including a primordial magnetic field we have seen that Einstein equations do not contain in its dynamics all the necessary information to determine the curvature tensor from empirical data. Namely, boundary conditions including the Weyl tensor, defined on a spatial hypersurface, cannot be provided by observational data, which lies in the past null cone of the observer. In other words, the Einstein equations should represent the universe as an open system rather than a closed totality, since at any moment new elements not included in the dynamical equations can play a role and modify the spacetime properties we observe. In particular, the presence of the Weyl tensor in this case shall change the perturbative version of the model and consequently the structure formation even if it cannot be perceived through the geodesic deviation.  The attempt in developing such cosmological model has also some features that should be explored further. For instance, the cases for which the tri-curvature is nonnull, which gives $k$ the status of a bifurcation parameter, should be addressed. Also, if we fix the tri-curvature to be flat but instead let $k$ assume the opposite sign we have considered in this paper, thus violating the cosmic censorship, then we get the correct Newtonian term in the effective potential for test particles at small values of the radial coordinate. For future developments, we also intend to investigate the relation between the magnetic field and the Weyl tensor interpreted as tidal forces from first principles using kinetic theory.  In summary, we have shown how the standard cosmological model sets the Weyl tensor equal to zero \\textit{ab initio} and that this is not a consequence of the Friedmann equations. This was shown through the development of a cosmological model with constant spatial curvature and a non-zero Weyl tensor without spoiling the conventional time evolution of the universe through the Friedmann equations."
    },
    "1402/1402.1548_arXiv.txt": {
        "abstract": "We describe the development of a noise-temperature testing capability for phased-array antennas operating in receive mode from 0.7~GHz to 1.8~GHz. Sampled voltages from each array port were recorded digitally as the zenith-pointing array under test was presented with three scenes: (1) a large microwave absorber at ambient temperature, (2) the unobstructed radio sky, and (3) broadband noise transmitted from a reference antenna centred over and pointed at the array under test.  The recorded voltages were processed in software to calculate the beam equivalent noise temperature for a maximum signal-to-noise ratio beam steered at the zenith.  We introduced the reference-antenna measurement to make noise measurements with reproducible, well-defined beams directed at the zenith and thereby at the centre of the absorber target.  We applied a detailed model of cosmic and atmospheric contributions to the radio sky emission that we used as a noise-temperature reference.  We also present a comprehensive analysis of measurement uncertainty including random and systematic effects.  The key systematic effect was due to uncertainty in the beamformed antenna pattern and how efficiently it illuminates the absorber load.  We achieved a combined uncertainty as low as 4~K for a 40~K measurement of beam equivalent noise temperature.  The measurement and analysis techniques described in this paper were pursued to support noise-performance verification of prototype phased-array feeds for the Australian Square Kilometre Array Pathfinder telescope. ",
        "introduction": "Developing low-noise, wideband, receive-only array antennas is crucial to delivering the Square Kilometre Array (SKA) telescope \\citep{Dewdney2009}.  Using aperture arrays and phased-array feeds (PAFs) allows more information to be collected from more of the sky in parallel.  This increases instantaneous field of view, increases survey speed, and allows more agile observing strategies as electronic beam steering can be immediate.  Array antennas enable telescope designers to spend more money on digital signal processing and less on mechanical signal processing via telescope dishes for a fixed performance goal.  This trade-off becomes more effective with time because digital signal-processing is becoming exponentially cheaper while the cost of dishes is not.  The SKA project explored this trade-off \\citep{Schilizzi2007,Chippendale2007,Alexander2007} and settled on significant deployments of both PAFs and aperture arrays in SKA phase 1 \\citep{Dewdney2013}.   \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth, angle=0]{DSCN1872_HotState_Leveled.eps} \\caption{Absorber rolled over array under test at Parkes.}\\label{fig:hot-load} \\end{center} \\end{figure}  Important to the development of low-noise array antennas is the ability to make accurate and reproducible measurements of their noise performance after beamforming.  A common approach for measuring array noise temperature is to apply the same Y-factor method used for single-antenna astronomy receivers \\citep{Sinclair1991} to the beamformed power from an array antenna \\citep{Woestenburg2003}.  The Y-factor is the ratio of beamformed power between observations of ``hot'' and ``cold'' loads.  At decimetre wavelengths, the hot load is often provided by microwave absorber at ambient temperature (Figure \\ref{fig:hot-load}) and the cold load by cosmic radio emission from the unobstructed sky.  A number of groups have reported recent developments in test facilities and measurement techniques for low-noise arrays.  Some experiments have positioned the absorber at the end of a tapered metal funnel or ground shield \\citep{Warnick2009,Woestenburg2011}. Others have used more open structures, like CSIRO's in Figure \\ref{fig:hot-load}, to support the absorber over the array under test \\citep{Woestenburg2012}.  Previous work with a ground shield has indicated small differences between measured system noise temperatures with and without the shield \\citep{Woestenburg2011}. The differences generally decrease with increase in the beamformed directivity of the array under test. At low directivity, where the shield is significant, the shield only partly decreased the effect of the terrestrial environment.  In this paper we describe CSIRO's development of an aperture-array noise-temperature testing capability at Parkes Observatory.   We develop a Y-factor approach similar to \\citep{Woestenburg2012} but introduce a reference-antenna (Figure \\ref{fig:absorber}) measurement to constrain the pointing of the beam towards the zenith and therefore the centre of the absorber. ",
        "conclusions": "We have demonstrated the measurement of partial beam equivalent noise temperature as low as ${\\hat{T}_\\text{n}=40 \\ \\text{K}}$ with a combined standard uncertainty (estimated standard deviation) as low as ${u_\\text{c}(\\hat{T}_\\text{n})=4 \\ \\text{K}}$.  This combined uncertainty was dominated by uncertainty $u(\\alpha)$ in the efficiency with which the beamformed array pattern illuminates the absorber.  The prioritised action-list for reducing uncertainty further is: \\begin{enumerate} \\item Reducing uncertainty in absorber-illumination efficiency $\\alpha$ by: \\begin{itemize} \\item increasing the solid angle subtended by the absorber by increasing its size or moving it closer to the array under test, \\item adding a ground shield to reflect more of the antenna pattern onto the absorber, or \\item accurately measuring or modelling the beamformed antenna pattern. \\end{itemize} \\item Moving the reference radiator into the far field of the array under test. \\item Increasing integration time for the ``cold'' sky measurement. \\item Reducing uncertainty in beam equivalent noise temperature due to radio emission from the sky $T_{\\text{ext},\\text{sky}(A)}$ by: \\begin{itemize} \\item using the coldest possible region of the sky for the ``cold'' sky measurement, \\item accurately measuring or modelling the beamformed antenna pattern, and \\item improving the accuracy of the global sky model. \\end{itemize}  \\end{enumerate} Addressing items (1) to (3) would reduce the median combined standard uncertainty to just ${u_\\text{c}(\\hat{T}_\\text{n})=2 \\ \\text{K}}$ over 0.7~GHz to 1.8~GHz."
    },
    "1402/1402.4471_arXiv.txt": {
        "abstract": "X-ray observations of circumgalactic coronae provide a valuable means by which to test galaxy formation theories. Two primary mechanisms are thought to be responsible for the establishment of such coronae: accretion of intergalactic gas and/or galactic feedback. In this paper, we first compare our \\emph{Chandra} sample of galactic coronae of 53 nearby highly-inclined disc galaxies to an analytical model considering only the accretion of intergalactic gas. We confirm the existing conclusion that this pure accretion model substantially over-predicts the coronal emission. We then select 30 field galaxies from our original sample, and correct their coronal luminosities to uniformly compare them to deep X-ray measurements of several massive disc galaxies from the literature, as well as to a comparable sample of simulated galaxies drawn from the Galaxies-Intergalactic Medium Interaction Calculation (GIMIC). These simulations explicitly model both accretion and supernovae feedback and yield galaxies that exhibit X-ray properties in broad agreement with our observational sample. However, notable and potentially instructive discrepancies exist between the slope and scatter of the $L_X-M_{200}$ and $L_X-{\\rm SFR}$ relations, highlighting some known shortcomings of GIMIC, for example, the absence of AGN feedback, and possibly the adoption of constant stellar feedback parameters. The simulated galaxies exhibit a tight correlation (with little scatter) between coronal luminosity and halo mass. Having inferred $M_{200}$ for our observational sample via the Tully-Fisher relation, we find a weaker and more scattered correlation. In the simulated and observed samples alike, massive non-starburst galaxies above a typical transition mass of $M_*\\sim2\\times10^{11}\\rm~M_\\odot$ or $M_{200}\\sim10^{13}\\rm~M_\\odot$ tend to have higher $L_X/M_*$ and $L_X/M_{200}$ than low-mass counterparts, indicating that the accretion of intergalactic gas plays an increasingly important role in establishing the observable hot circumgalactic medium with increasing galaxy mass. ",
        "introduction": "\\label{PaperIIIsec:Introduction}  The presence of extended X-ray emitting coronae associated with present-day massive disc galaxies is a generic prediction of cold dark matter (CDM) galaxy formation theory (e.g., \\citealt{White91,Benson10}). The coronae serve as a reservoir of both diffuse, metal-poor gas accreted from the intergalactic medium (IGM), and metal-rich gas that is either ejected from galaxies by energetic feedback or stripped from infalling satellites. Two mechanisms are thought to be chiefly responsible for establishing the the X-ray emission properties of coronae: feedback injected by stellar wind/supernovae (SNe) (e.g., \\citealt{Strickland00a}) and active galactic nuclei (AGN), and the shock heating and adiabatic compression of gas accreted from the IGM (e.g., \\citealt{White78}).  Although galactic coronae have been detected around many disc galaxies, their origin remains uncertain. For example, the coronal luminosity is observed to correlate well with many tracers of feedback activity [e.g., the infrared (IR) or radio luminosities, \\citealt{Strickland04b,Grimes05,Tullmann06b,Li08,Li13b}], and the heavy element abundance of the X-ray luminous gas is also indicative of enrichment by SNe (e.g., \\citealt{Martin02,Ji09,Li09,Bregman13,Li13b}). These are often regarded as direct evidence for the feedback scenario (but see \\citealt{Crain13} for an alternative interpretation). Furthermore, non-gravitational heating is also indicated by the change of slope of the scaling relations between the coronal luminosity/temperature and the galaxy mass in different mass ranges (e.g., \\citealt{Ponman99,Ponman03,OSullivan03}). The uncertainty is driven by the difficulty of detecting and characterizing the diffuse thermal X-ray emission at large galactocentric radii (e.g., \\citealt{Benson00,Rasmussen09}). In recent years, with either stacking of large numbers of X-ray observations of various types of galaxies \\citep{Anderson13} or moderately deep single-pointing X-ray observations of massive disc galaxies \\citep{Anderson11,Dai12,Bogdan13a,Bogdan13b}, X-ray emission from the putative coronae on large scales (at radii up to $\\sim50\\rm~kpc$) has been reported around several galaxies which are not particularly active in SF. These coronae are considered to be most likely produced by the accretion of intergalactic gas. However, such detections are still rare and only concern galaxies within a narrow range of stellar mass (see Table~\\ref{table:GalaxyPara}). Comparison with X-ray measurements of galaxies in a broad mass range is thus particularly desirable.  In our previous studies \\citep{Li13a,Li13b} (hereafter, Paper~I, II), we investigated the galactic coronae associated with a sample of \\emph{Chandra}-observed nearby highly-inclined disc galaxies. We found a tight correlation between the coronal luminosity and the SN energy input rate; the abundance ratio of the coronal gas is also consistent with a combined contribution from core collapsed and Type~Ia SNe. These results can be qualitatively interpreted, naively, as evidence that galactic coronae are mainly established by SN feedback. However, our sample only includes low and intermediate mass galaxies (typically with stellar mass $M_*\\lesssim 2\\times10^{11}\\rm~M_\\odot$), which may not be massive enough for accreted gas to contribute a dominant fraction of the detected X-ray flux. In fact, recent cosmological hydrodynamical simulations invoking both accretion and feedback predict a coronal luminosity range that is broadly consistent with X-ray observations \\citep{Crain10a}. These simulations trace the dynamics of the coronal gas, a large fraction of which is quasi-hydrostatic or even infalling. \\citet{Crain13} further showed that the high ($\\sim$solar) metallicities inferred from the X-ray spectroscopy are not in conflict with models in which the bulk (by mass) of the hot circumgalactic medium (CGM) is established by accretion. The soft X-ray emission is dominated by collisionally-excited metal ions deposited by feedback, whilst the majority of the accreted gas is metal poor and is comparatively radiatively inefficient. The hot CGM thus appears to be typically metal-poor in a mass-weighted sense, but its X-ray luminosity-weighted metallicity is often close to solar. Therefore, accretion of intergalactic gas may still be important, especially in those massive galaxies not included in our \\emph{Chandra} sample. In Paper~II, we also compared our measurements to X-ray observations of massive elliptical galaxies, which have significantly steeper $L_X-M_*$ relations. The formation of such galaxies may, however, be subject to more violent mergers and the corresponding starbursting process, potentially complicating the interpretation of observations of their coronae. Furthermore, the typically rich clustered environment of massive elliptical galaxies also potentially contaminates the measurement of coronal properties associated with individual galaxies. Therefore, comparisons with X-ray measurements of elliptical galaxies do not offer the cleanest means by which to place direct constraint on disc galaxy formation models.  We herein conduct a quantitative comparison of our X-ray measurements of disc galaxies to the X-ray measurements of several isolated massive disc galaxies from the literature (\\S\\ref{PaperIIIsubsec:MassiveGalaxy}), as well as the results of hydrodynamical simulations invoking both accretion and feedback (i.e., the GIMIC simulations, see \\S\\ref{PaperIIIsubsec:Simulations}), in order to further explore the origin of coronae around disc galaxies and place direct constraint on galaxy formation models. This paper is organized as follows. In \\S\\ref{PaperIIIsec:data}, we describe our procedure for correcting coronal luminosity measurements to ensure data homogenization between the galaxies selected from our original \\emph{Chandra} sample, galaxies from the literature, and the GIMIC simulations. We compare the observational results to an analytical model as well as the GIMIC simulations in \\S\\ref{PaperIIIsec:Comparison}, and further discuss the scientific implications in \\S\\ref{PaperIIIsec:Discussion}. Our results and conclusions are summarized in \\S\\ref{PaperIIIsec:Summary}. ",
        "conclusions": "\\label{PaperIIIsec:Discussion}  \\subsection{Effects of Missing AGN Feedback}\\label{PaperIIIsubsec:AGNfeedback}  As introduced in \\S\\ref{PaperIIIsubsec:Simulations}, the GIMIC simulations do not include AGN feedback. The absence of AGN feedback in GIMIC may result in the over-prediction of the number density of \\emph{massive} disc galaxies \\citep{Crain09,Crain10a}. The detailed operation of AGN feedback on galaxy-wide scales remains ill-understood, but is expected to be important in galaxies with massive spheroids, where the energy output could heat and eject gas, reducing the emissivity in the central regions (e.g., \\citealt{McCarthy10,McNamara12} and references therein). We have adopted a high-mass cut of $M_*\\sim10^{11.7}\\rm~M_\\odot$ to the GIMIC galaxies to minimize the effect of missing AGN feedback (\\S\\ref{PaperIIIsubsec:Simulations}). Generally, this AGN effect is less important than Type~Ia SNe feedback for galaxies with $M_*\\lesssim10^{11}\\rm~M_\\odot$ (e.g., \\citealt{David06}), and should not qualitatively affect the comparison between observation and simulation in the mass range of our \\emph{Chandra} sample. For galaxies with $M_*\\gtrsim10^{11}\\rm~M_\\odot$, however, as shown in Fig.~\\ref{fig:M200Scaling}a, the stellar masses at fixed $M_{200}$ are inconsistent with estimates derived from observational techniques \\citep{Leauthaud12a}. This over-prediction of $M_*$ is most likely a result of the enhanced radiative cooling and SF in these massive galaxies due to the absence of AGN feedback in GIMIC. Correspondingly, this missing AGN feedback may also explain the high luminosity of coronal X-ray emission at a given $M_{200}$, for $M_{200}\\gtrsim a~few\\times10^{12}\\rm~M_\\odot$ (Fig.~\\ref{fig:compareGIMIC}b).  \\begin{figure}[!h] \\begin{center} \\epsfig{figure=fig03.eps,width=1.0\\textwidth,angle=0, clip=} \\caption{(a) $M_{200}$ v.s. $M_*$. (b) $M_{200}$ v.s. $v_{rot}$. Symbols are the same as those in Fig.~\\ref{fig:compareGIMIC}. The solid curve in panel (a) is the SMHMR from \\citet{Leauthaud12a} [for the model $SIG\\_MOD2$ (with the modeling of the stellar mass measurement errors) and the lowest redshift bin ($z=0.22-0.48$)]. The solid line in panel (b) marks the $M_{200}-v_{rot}$ relation of \\citet{Navarro97}, which is used to compute $M_{200}$ of most of the observed galaxies, except for NGC~266.}\\label{fig:M200Scaling} \\end{center} \\end{figure}  To first order, the rotation velocity of galaxies is determined by the mass and concentration of their parent halo. Whilst the concentration of halos modeled in collisionless simulations is known to follow a tight relation (e.g., \\citealt{Navarro97,Bullock01}), the co-evolution of baryons and dark matter in more realistic simulations induces systematic deviations and scatter in this relation (often described as a ``back reaction'', e.g., \\citealt{Duffy10}). As shown in Fig.~\\ref{fig:M200Scaling}b, low-mass galaxies ($M_{200}\\lesssim (2-3)\\times10^{12}\\rm~M_\\odot$) from GIMIC exhibit a larger scatter in $v_{rot}$ than more massive counterparts. The low-$v_{rot}$ boundary of low-mass galaxies is roughly consistent with \\citet{Navarro97}'s relation (derived from collisionless simulations), and is used here to estimate $M_{200}$ of all the observed galaxies except for NGC~266. On the other hand, the high-$v_{rot}$ boundary of low-mass galaxies seems consistent with high-mass ones. In contrast, high-mass galaxies with $M_{200}\\gtrsim (2-3)\\times10^{12}\\rm~M_\\odot$ have clearly lower scatter in $v_{rot}$, and no galaxies have $v_{rot}$ as low as predicted by \\citet{Navarro97}'s relation. In the absence of AGN feedback, the halo of a massive galaxy will tend to over-cool and over-contract, resulting in a greater rotation velocity. Therefore, the absence of AGN feedback in GIMIC, particularly for high-mass galaxies, likely results in an artificially high baryonic matter density at small radii, and thus a higher dark matter density and maximum rotation velocity.  \\subsection{Role of SF in the Coronal X-ray Emission}\\label{PaperIIIsubsec:CoolHotGasInteraction}  Our sample is selected based on the available \\emph{Chandra} data. Archival observations are often proposed for active starburst galaxies which tend to have luminous X-ray halos. Among the 30 galaxies used in the present paper, $\\sim40\\%$ are classified as ``starburst'' ($\\sim35\\%$ if one includes the four massive disc galaxies; Table~\\ref{table:GalaxyPara}), while in the GIMIC \\textit{sample} (not necessarily the simulation taken as a whole), this fraction is only $\\sim9\\%$ (defining a simulated galaxy with ${\\rm SFR}/M_*\\gtrsim0.1\\rm~Gyr^{-1}$ as a starburst, broadly consistent with the definition for the observed galaxies; Fig.~\\ref{fig:SFproperties}a). This low fraction in the sample is likely a consequence of the exclusion from the sample of interacting galaxies. We thus separated starburst and non-starburst galaxies in the above comparisons (\\S\\ref{PaperIIIsubsec:compareGIMIC}).  The coronal luminosity discrepancy between the observed starburst and the GIMIC galaxies may be partly attributed to this difference in the sample selection. However, at small radii, the X-ray intensity always drops for GIMIC galaxies \\citep{Crain10a,Crain13}, but increases for the observed ones (Paper~I). As presented in \\S\\ref{PaperIIIsubsec:compareGIMIC}, the coronal luminosity of our \\emph{Chandra} sample galaxies in $0.01-0.1r_{200}$ is on average $\\sim15\\%$ of those adopted in Papers~I and II (\\S\\ref{PaperIIIsubsec:compareGIMIC}), while the corrected coronal luminosity of the GIMIC galaxies in $0.01-0.1r_{200}$ is typically $\\sim(20-50)\\%$ of the luminosity in the entire halo (\\S\\ref{PaperIIIsubsec:Simulations}). Considering that the X-ray luminosity adopted in Papers~I and II is also measured in the inner region, we could conclude that GIMIC under-predict the coronal X-ray emission of low-mass galaxies at least at small radii, although the chosen of photometry region of $r=0.01-0.1r_{200}$ has minimized this discrepancy. In fact, this under-prediction may also exist at $0.01-0.1r_{200}$, as indicated by the steeper slope of the $L_X-{\\rm SFR}$ relation for the GIMIC galaxies (with respect to the observed galaxies in our \\emph{Chandra} sample; Fig.~\\ref{fig:compareGIMIC}d). We consider the SFR range of $0.5-5\\rm~M_\\odot~yr^{-1}$, because galaxies with larger SFRs are likely massive ones in GIMIC and starburst in our \\emph{Chandra} sample, while those with a lower SFR in GIMIC have too large scatter in $L_X$ and many of them are indeed massive galaxies (see below; Fig.~\\ref{fig:compareGIMIC}d). We find an average coronal luminosity of $\\log (L_X/{\\rm ergs~s^{-1}})=38.89_{-0.16}^{+0.15}$ for our observed galaxies, compared to $\\log (L_X/{\\rm ergs~s^{-1}})=38.02_{-0.10}^{+0.09}$ for the simulated ones.  \\begin{figure}[!h] \\begin{center} \\epsfig{figure=fig04.eps,width=1.0\\textwidth,angle=0, clip=} \\caption{SF properties of the sample galaxies. (a) $L_X$ v.s. the specific SFR per unit stellar mass (${\\rm SFR}/M_*$). (b) SFR v.s. $M_*$. (c) SFR v.s. $M_{200}$. Symbols are the same as those in Fig.~\\ref{fig:compareGIMIC}.}\\label{fig:SFproperties} \\end{center} \\end{figure}  The most intuitive reason for this under-prediction of X-ray luminosity around $L^\\star$ non-starburst field galaxies may be from the adopted physical and/or numerical ingredients in GIMIC. Many physical assumptions and parameters, as well as the adopted numerical methods (e.g., \\citealt{Scannapieco12}), may affect the simulated galaxy properties, e.g., the adopted feedback implementation, metal line cooling, stellar initial mass function, and the assumed cosmology (e.g., \\citealt{Haas13a,Haas13b}). In particular, the choice of feedback parameters plays a key role in reproducing the observed SF properties of galaxies (e.g., \\citealt{Schaye10,Scannapieco12,Vogelsberger13}), and possibly the related halo gas cooling and X-ray emission. As suggested in many cosmological simulations, there exists a fundamental problem that the specific SFRs of galaxies (particularly low-mass galaxies) at low-redshift are significantly under-predicted (e.g., \\citealt{Weinmann12}). It has been suggested that this discrepancy arises because both semi-analytic and hydrodynamical simulations couple the growth of galaxies too strongly to the growth of their host dark matter haloes. Since haloes assemble early, simulated galaxies do so too, and thus form significantly earlier than is observed in nature (e.g., \\citealt{Weinmann12}). Similarly, GIMIC also under-predicts the SFR of low-mass galaxies \\citep{Crain09}. This under-prediction of current SF activity may have potential effects on the X-ray emission, e.g., via related stellar feedback and cool-hot gas interaction, but the net effect is not yet well understood. For example, there exists a small number of low-SFR high-$L_X$ galaxies in the GIMIC simulations (Fig.~\\ref{fig:compareGIMIC}d). As indicated in Fig.~\\ref{fig:SFproperties}b,c, these galaxies are most likely massive ones, which are apparently post-starbursts whose SF have already ceased but the radiative cooling of the coronal gas is still strong. Some observed galaxies (e.g., NGC~2787, NGC~2841, NGC~4594, and NGC~5170) may have similar SF properties, i.e., with low SFR but high $M_*$. The SF activities in these galaxies have probably been quenched, e.g., via morphological quenching or other mechanisms (e.g., \\citealt{Martig09,Li09,Li11} and references therein). However, the diffuse X-ray luminosities of these galaxies are much lower than those of other massive galaxies with higher SFR (e.g., the massive disc galaxies from the literature; Fig.~\\ref{fig:compareGIMIC}), apparently inconsistent with what is predicted by GIMIC.  The above comparisons clearly demonstrate that X-ray observations can place important constraints on the galaxy formation theory, especially its feedback implementation. Improved feedback prescriptions, especially in low-mass galaxies, may need to be considered in next generation simulations to correctly reproduce both the current SFR and the coronal X-ray luminosity. In particular, GIMIC adopted constant kinetic feedback parameters (initial wind velocity of $600\\rm~km~s^{-1}$ and mass loading factor of $\\dot{m}_{wind}/\\dot{m}_*=4$), which were chosen to scale the global SFR density to correspond to observational estimates. Although the adopted efficiency (the wind energy accounts for $\\sim80\\%$ of the SN energy) and mass loading factor are roughly in agreement with observational estimates [e.g., for M82 by \\citet{Strickland09}], it is likely that the macroscopic properties of outflows correlate with the detailed structure of the ISM (e.g., \\citealt{Efstathiou00,Creasey13}) and, by extension, the properties of galaxies. The simplicity of the adopted physical and/or numerical ingredients in GIMIC is a possible reason that the SF and X-ray properties of the simulated galaxies are too strongly coupled with the dark matter halo mass of the galaxy (e.g., Fig.~\\ref{fig:compareGIMIC}b). It will be interesting to compare the X-ray scaling relations of GIMIC galaxies with those of galaxies formed in next-generation hydrodynamical simulations adopting feedback schemes calibrated to reproduce the stellar mass function of local galaxies.  Another potential explanation for the under-prediction of soft X-ray emission is the adoption of a single-phase ISM in GIMIC. The gas is assigned an average temperature within the numerical resolution, which is typically insufficient to resolve a single massive SF region. In these SF regions, cold gas often dominates the total gas mass, while hot gas dominates the X-ray emission. Therefore, averaging the thermal state of the gases will result in a too low temperature to emit X-ray efficiently, thus significantly under-predict the X-ray emission. This effect is the most significant near a galactic disc, where the cool/hot gases strongly mix with each other, and the X-ray emission can be enhanced via enhanced radiative cooling of mass loaded cool gas and charge exchange (e.g., \\citealt{Li11,Liu11,Liu12}). This effect may also help to explain the declining X-ray intensity profile toward small radii in GIMIC \\citep{Crain10a,Crain13}, which is not typically observed in real galaxies.  As a general conclusion, GIMIC likely under-predict the role of SF in X-ray emission for low mass galaxies. There are two origins of this under-prediction: the under-prediction of the SFR for low-mass galaxies due to the too strong coupling of the growth of galaxies to the growth of their host dark matter halos, and the under-prediction of X-ray emission from SF due to the assumption of the single-phase ISM.  \\subsection{Differences between Low and High Mass Galaxies}\\label{PaperIIIsubsec:DifferenceLowHighMass}  Although the $L_X-M_{200}$ or $L_X-M_*$ slopes themselves are not well constrained due to the large scatter of the observational data, it is clear that the four massive disc galaxies (NGC~266, NGC~1961, NGC~6753, UGC~12591) are more X-ray luminous than the non-starburst galaxies below a typical transition mass of $M_*\\sim2\\times10^{11}\\rm~M_\\odot$ or $M_{200}\\sim10^{13}\\rm~M_\\odot$ (Fig.~\\ref{fig:compareGIMIC}b,c). As argued in previous studies (e.g., \\citealt{Bogdan13a,Bogdan13b}), this high X-ray luminosity is most likely a signature of the accretion of intergalactic gas.  There are several possible mechanisms which may produce the high X-ray luminosity of these massive disc galaxies. \\emph{Firstly}, the deeper gravitational potential, and possibly the stronger surrounding thermal and ram-pressure (e.g., \\citealt{DallaVecchia08,Lu11}), may help the galaxy to retain more hot gas within its halo, resulting in an increasing baryon fraction (especially the hot gas fraction) with the halo mass (e.g., \\citealt{Crain10a,Dai12}). \\emph{Secondly}, the change of the $L_X/M_*$ ratio could also be a result of the change of the accretion mode. Numerical simulations have confirmed the change of the cold to hot accretion from below to above the (baryonic) mass of $\\sim10^{10.3}\\rm~M_\\odot$ of a galaxy (e.g., \\citealt{Keres05,Keres09,Crain10a,vandeVoort11}). Above this transition mass, a larger fraction of gas can be gravitationally heated to an X-ray emitting temperature in the galaxy vicinity. \\emph{Finally}, the dynamical state of the coronal gas may also affect its overall X-ray emissivity by changing its radial distribution \\citep{Ciotti91,OSullivan03}. A massive galaxy with a corona that tends to stay in a hydrostatic or even inflow state can have a steeper density profile (higher gas density in the center), enhancing the X-ray emission. Conversely, a galaxy with its corona in a subsonic outflow state tends to have a flatter density profile and may be less luminous in the galactic inner regions (e.g., \\citealt{Tang09a}). In all cases, the dark matter halo mass appears to be a key parameter in determining the coronal luminosity and it is likely that the accretion of intergalactic gas plays an increasingly important role in producing the galactic coronae above the transition mass.    In order to study the coronae around disc galaxies, we have constructed a \\emph{Chandra} database of 53 nearby highly-inclined disc galaxies (Paper~I), and have conducted a correlation analysis of their coronal and other multi-wavelength properties (Paper~II). In this paper, we have compared our results to the predictions from a simple analytical model considering only the accretion of intergalactic gas (in the absence of feedback), the measurements of several massive disc galaxies from the literature, as well as results from recent cosmological hydrodynamical simulations invoking both accretion and feedback (the GIMIC simulations). Our main results and conclusions are set out below.  The observed X-ray emission in the vicinity of galactic discs (those explored in our \\emph{Chandra} observations without further corrections in \\S\\ref{PaperIIIsubsec:OurSample}) is much less ($\\sim2\\%$) than predicted by a simple analytical model considering only the accretion of intergalactic gas. Furthermore, the radiative cooling rate of coronal gas in the same radial aperture is typically too low ($\\lesssim10\\%$ of the SFR) to be the primary source of gas to maintain the on-going SFR in the galaxies. The GIMIC simulations more accurately reproduce the observed X-ray scaling relations. They broadly reproduce the luminosity range of the coronal X-ray emission in $0.01-0.1r_{200}$ of $L^\\star$ galaxies, including the scatter in $L_X$ at a given $v_{rot}$ or $M_*$. However, the $L_X-M_{200}$ relation of simulated galaxies differs from that inferred from observations, both in terms of the slope and scatter.  For the observed galaxies, low-mass starbursts appear to be more X-ray luminous than more quiescent counterparts at a given galaxy mass. There is a tight $L_X-{\\rm SFR}$ correlation, even for corona at large radii of $r=0.01-0.1r_{200}$, similar to that of the X-ray emission at smaller radii. The overall trends of the observed and simulated galaxies on the $L_X-{\\rm SFR}$ plot show little similarity, indicating that GIMIC potentially under-predicts the role of SF in producing the X-ray emission. There are two possible origins of this under-prediction: the under-prediction of the SFR for low-mass galaxies due to the too strong coupling of the growth of galaxies to the growth of their host dark matter halos, and the under-prediction of X-ray emission from SF due to the assumption of the single-phase ISM.  Coronal X-ray luminosity increases much faster with galaxy mass above a typical transition mass of $M_*\\sim2\\times10^{11}\\rm~M_\\odot$ or $M_{200}\\sim10^{13}\\rm~M_\\odot$. Below this mass, the $L_X-M_{200}$ and $L_X-M_*$ relations of the observed galaxies can both be well characterized with a linear scaling relation. The higher $L_X/M_*$ or $L_X/M_{200}$ ratio of galaxies above the transition mass indicates that a massive disc galaxy tends to have a massive dark matter halo, to be dominated by hot-mode accretion, and/or to host a corona most likely in a hydrostatic state. The accretion of intergalactic gas likely plays an increasingly important role in producing the galactic coronae with increasing galaxy mass.  The above results demonstrate that X-ray observations can place important constraints on the galaxy formation theory, especially the astrophysics of accretion and feedback. Although the simulations have made a great progress to reproduce the observed X-ray properties of galaxies, some significant discrepancies still exist and are quite suggestive. The absence of AGN feedback in GIMIC tends to over-predict the number of massive galaxies, as well as a simultaneous over-prediction of $v_{rot}$, $M_*$, and $L_X$ at a given $M_{200}$ for massive galaxies. In addition, GIMIC may yield too strong a coupling of the galaxy and coronal properties ($v_{rot}$, $M_*$, SFR, $L_X$) to the host dark matter halo mass, which is one of the major origins of the discrepancies between the observations and simulations. Furthermore, the adoption of constant feedback parameters (mass loading and injection velocity) is likely too simplistic. It is also important to study the X-ray properties of low-mass galaxies, which are not massive enough to retain hot coronae, and the X-ray emission is mostly from stellar feedback and the related cool-hot gas interaction, which is poorly understood from current numerical simulations.  The present paper is based on the archival observations and X-ray measurements from the literature. Most of these observations are too shallow to probe the faint X-ray emission from the outskirts of the galactic halos. In addition, the sample is still not large nor uniform. In particular, very few galaxies are observed in X-rays around or above the transition mass. We have also focused on the \\emph{X-ray luminosities} of the galactic coronae. Most of the archival observations are too shallow to yield valuable constraints on hot gas abundance, and hence on gas density and total baryon content of the galaxies (e.g., \\citealt{Bregman07}). Therefore, deeper X-ray observations of more galaxies spanning the entire galaxy mass range are highly desirable in order to put tighter constraints on the galaxy formation models and further improve our understanding of the global properties of galactic coronae."
    },
    "1402/1402.6315.txt": {
        "abstract": "We combine data from the Spitzer Survey for Stellar Structure in Galaxies (S$^4$G), a recently calibrated empirical stellar mass estimator from Eskew et al., and an extensive database of H{\\small I} spectral line profiles to examine the baryonic Tully-Fisher (BTF) relation. We find 1) that the BTF has lower scatter than the classic Tully-Fisher (TF) relation and is better described as a linear relationship, confirming similar previous results, 2) that the inclusion of a radial scale in the BTF decreases the scatter but only modestly, as seen previously for the TF relation,  and 3) that the slope of the BTF, which we find to be $3.5\\pm0.2$ ($\\Delta \\log M_{baryon}/\\Delta \\log v_c$), implies that on average a nearly constant fraction ($\\sim 0.4$) of all baryons expected to be in a halo are ``condensed\" onto the central region of rotationally supported galaxies. The condensed baryon fraction, $M_{baryon}/M_{total}$, is, to our measurement precision, nearly independent of galaxy circular velocity (our sample spans circular velocities, $v_c$, between 60 and 250 km s$^{-1}$, but is extended to $v_c\\sim 10$ km s$^{-1}$ using data from the literature). The observed galaxy-to-galaxy scatter in this fraction is generally $\\le$ a factor of 2 despite fairly liberal selection criteria. These results imply that cooling and heating processes, such as cold vs. hot accretion, mass loss due to stellar winds, and AGN driven feedback, to the degree that they affect the global galactic properties involved in the BTF, are independent of halo mass for galaxies with $10 < v_c < 250$ km s$^{-1}$ and typically introduce no more than a factor of two range in the resulting $M_{baryon}/M_{total}$. Recent simulations by Aumer et al. of a small sample of disk galaxies are in excellent agreement with our data, suggesting that current simulations are capable of reproducing the global properties of individual disk galaxies. More detailed comparison to models using the BTF holds great promise, but awaits improved determinations of the stellar masses. ",
        "introduction": "\\label{sec:intro}  Empirical galaxy scaling relations are a testament to the existence of underlying physical principles of galaxy formation. Among such scaling relations, the relationship between the width of the neutral hydrogen line and the luminosity of a galaxy \\citep{tf} stands as one of the most useful \\citep[for example see its use as a distance estimator;][]{freedman} and constraining \\citep[for example see its use to test complicated baryonic physics in dark matter halos;][]{steinmetz}.  Even though the Tully-Fisher relationship (hereafter the TF relation) has been extensively vetted and explored \\citep{freedman10}, some questions remained unanswered. Foremost are those relating to the physical origin of the relationship. Although vaguely related to the Virial theorem, the relationship is not simply a recasting of that theorem \\citep[cf.][]{mcgaugh98,mcgaugh00,z12a}. In particular, one can imagine constructing two galaxies with the same rotation curve, but extending the stellar disk a few times farther out in one. Although both galaxies would satisfy the Virial theorem, the two galaxies could not lie on the same relationship between rotation velocity and luminosity. Nature, apparently, cannot envision two such galaxies.  However, the TF relation in its simplest incarnation is not a complete description of all disk galaxies. It fails to match the characteristics of some faint, gas rich galaxies \\citep{carignan,persic,meurer,mcgaugh00}. A scaling relationship is recovered if one recasts it as one between rotation velocity and baryonic mass rather than just luminosity \\citep{freeman,walker,mcgaugh00,verheijen,geha}, and this is referred to as the baryonic TF (hereafter, BTF).  These results suggest that the original TF exists because for most galaxies in TF studies the luminosity is a reasonably precise proxy for stellar mass and the gas mass is negligible.  In a quest to uncover more clues and improve the fidelity of the TF as a distance estimator, investigators have long sought additional parameters that would help reduce the scatter in the TF relation --- a search for a second, or even third, parameter beyond the rotational velocity \\citep[see, for examples,][]{strauss,zwaan,sprayberry,courteau97,courteau,mcgaugh05b,kassin,hall}. In this context, evidence for the importance of a scaling radius has been presented \\citep[cf.][]{kassin,hall}, although many other studies found no such dependence \\citep{zwaan, sprayberry, courteau, mcgaugh05b}.  The principal scaling relationship for early type galaxies, the Fundamental Plane \\citep{dd87,d87}, includes a kinematic term, a scaling radius, and a surface brightness term. Following along that line of reasoning, we have explored a joint scaling relationship for both late and early types that has the flavor of the Fundamental Plane  \\citep{z06a,z06b,z08,z12a}. Given its antecedent and geometry, it is referred to as the Fundamental Manifold. The implication of that work in the current context is that disk galaxies should show residuals from the standard TF that correlate with the half light radius. In this study, we examine whether the introduction of this scaling radius improves the BTF as well.  Finally, the BTF has great promise to provide detailed, quantitative tests of galaxy formation and evolution models \\citep[for some examples of this approach see:][]{mayer,gnedin,governato,avila,gurovich,mcgaugh12,aumer}. The limiting uncertainty in such tests has been the stellar mass determination. While prescriptions for the mass-to-light ratio based on observed colors are available \\citep{bell}, these depend sensitively, at a level of precision that incapacitates the envisioned test, on the adopted stellar initial mass function and star formation history. Because these techniques are based on stellar population modeling they also carry fundamental uncertainties on the modeling of rare, but luminous, phases of stellar evolution \\citep{maraston,conroy}.  As an alternative, \\cite{eskew} used the resolved stellar population study of the Large Magellanic Cloud \\citep{harris} and {\\sl Spitzer} images \\citep{meixner} to calibrate the conversion of 3.6 and 4.5 $\\mu$m luminosities to stellar mass. This approach bypasses some, but not all, of the weaknesses mentioned above \\citep[see][for a discussion of the relative merits]{eskew}.  The use of the IR tracers also mitigates the role of internal extinction and detailed work at these wavelengths aims to correct for dust emission as well \\citep{meidt}. With the advent of this new stellar mass estimator and large samples of galaxies observed with the {\\sl Spitzer} Space Telescope, we now return to reexamine the BTF.  In this study we combine homogeneous, high quality IR data (3.6$\\mu$m and 4.5$\\mu$m surface photometry) from the {\\sl Spitzer} Survey for Stellar Structure in Galaxies \\citep[S$^4$G;][]{sheth}, which is less susceptible to extinction than optical data and has a more uniform stellar mass-to-light ratio, the \\cite{eskew} stellar mass estimator, and  H{\\small I} spectral line profiles from the literature curated by the Cosmic Flows project \\citep{courtois11} to re-explore some of the questions raised by both the standard and baryonic Tully-Fisher relationship. A strength of this study is that the parent sample is fairly broadly selected to be a magnitude limited, volume limited sample. As such, the results are representative of galaxies in general rather than of a ``pristine\" sample intended to provide the tightest scaling relation or most reliable distances. In \\S2 we describe the data, present our findings regarding the BTF and the role of other parameters in \\S3, and conclude in \\S4. ",
        "conclusions": "We examine the Tully-Fisher (TF) and Baryonic Tully-Fisher (BTF) using the new sample of IR photometry provided by the S$^4$G survey \\citep{sheth,munoz}. The use of this large, homogenous sample brings the natural advantages of uniformity and large statistics, and also provides photometry that is less susceptible to reddening and stellar population variations. We apply a new empirical stellar mass estimator from \\cite{eskew} to obtain stellar masses for the sample. When combined with the available large, homogenized archival H{\\small I} data of the Cosmic Flows program and complemented with additional observations \\citep{courtois11}, we have a primary sample of 903 galaxies for study.  We confirm the superiority of the BTF for lower mass galaxies and that the overall scaling relation becomes noticeably more linear than the TF. Furthermore, we quantify the scatter introduced by various  lax selection criteria related to the inclination limits of the sample, the morphological type range of the sample, and the effect of peculiar velocities on nearby galaxies. In the end, we are able to identify a BTF with a scatter of 0.18 dex, without detailed pruning of galaxies with either morphological or kinematic anomalies. This is a fair sample to compare to simulations where such cuts are difficult to reproduce.  We demonstrate that physical scale, here parameterized by the half light radius, correlates with the residual in the BTF. However, the gain in precision is exceedingly slight and the standard BTF ($\\log M_{baryon}$ vs. $\\log v_c$) is to be preferred for this galaxy sample for its simplicity.  We measure a BTF slope of $3.5 \\pm 0.2$, consistent with the expectation where galaxies concentrate a fixed fraction of their initial baryons in the central, detectable components (stars and gas). Using the MW to normalize the relation, we show that the expectation of this simple model, where 0.07 of the total mass of the halo (or alternatively about 40\\% of the baryons) is ``condensed\" onto the central regions of the halo, is an excellent fit to the data. This agreement demonstrates that, independent of circular velocities for the range of velocities explored here, the resulting effects of physical processes that could have affected the global properties of these galaxies (outflows, inflows) are independent of halo mass. Our sample also supports galaxy-to-galaxy scatter in the condensed baryon fraction up to a factor of two. Untangling the observational and intrinsic scatter is difficult with the current data, but manifestly a next avenue that needs to be explored.  Our value of the BTF slope agrees with certain previous studies \\citep{avila, hall} and we present one possible explanation for the disagreement with other studies that found steeper slopes \\citep{mcgaugh12}. We suggest that the disagreement comes about from the normalization of the stellar mass estimation. While we do not have direct evidence favoring one mass normalization over another, we prefer ours because it results in closer agreement in the derived BTF relations when using independent samples of gas and star dominated galaxies. We show that combining the gas and star dominated galaxy samples, when different stellar mass normalizations are used, can result in slope differences consistent with what is found among the various studies. Our data also result in excellent agreement with the recent simulations of \\cite{aumer}. Although that agreement cannot be used to support our stellar mass estimates, it does suggest that simulations are approaching a level of sophistication rivaling the observational uncertainties in determining the internal structural properties of individual galaxies. Finally, and perhaps eventually most illuminating when the origin is understood, we note the continuing tension between the BTF results, which suggests a constant $M_{baryon}/M_{halo}$, and those from statistical studies of galaxies, such as abundance matching models, that find a decreasing $M_{baryon}/M_{halo}$ with decreasing $M_{halo}$."
    },
    "1402/1402.5850_arXiv.txt": {
        "abstract": "Super-Earths transiting nearby bright stars are key objects that simultaneously allow for accurate measurements of both their mass and radius, providing essential constraints on their internal composition. We present here the confirmation, based on \\textit{Spitzer} transit observations, that the super-Earth HD\\,97658\\,b transits its host star. HD\\,97658 is a low-mass ($M_*=0.77\\pm0.05\\,M_{\\odot}$) K1 dwarf, as determined from the \\textit{Hipparcos} parallax and stellar evolution modeling. To constrain the planet parameters, we carry out Bayesian global analyses of Keck-HIRES radial velocities, and \\textit{MOST} and \\textit{Spitzer} photometry.  HD\\,97658\\,b is a massive ($M_P=7.55^{+0.83}_{-0.79} M_{\\oplus}$) and large ($R_{P} = 2.247^{+0.098}_{-0.095} R_{\\oplus}$ at 4.5 $\\mu$m) super-Earth. We investigate the possible internal compositions for HD\\,97658\\,b. Our results indicate a large rocky component, by at least 60\\% by mass, and very little H-He components, at most 2\\% by mass. We also discuss how future asteroseismic observations can improve the knowledge of the HD\\,97658 system, in particular by constraining its age. Orbiting a bright host star, HD\\,97658\\,b will be a key target for coming space missions TESS, CHEOPS, PLATO, and also JWST, to characterize thoroughly its structure and atmosphere. ",
        "introduction": "Among the diversity of the exoplanetary population, the class of the so-called ``super-Earths\", with a mass of a few Earth masses, is of utmost interest. Firstly, they do not exist in the Solar System, leaving us with only extrasolar planets for detailed studies. Secondly, they seem to be extremely common in the Galaxy \\citep{2013ApJ...766...81F,2013ApJ...770...69P,2013Sci...340..572H}. Thirdly, a wide variety of interiors are theoretically possible for a small range of radii ($1 < R_{\\rm P}/R_{\\oplus} < 5$) and masses ($1 < M_{\\rm P}/M_{\\oplus} < 10$), from pure iron composition with no atmosphere, to water planets with a significant atmosphere \\citep{2007ApJ...665.1413V,2007ApJ...669.1279S}. Transiting super-Earths around bright nearby stars are therefore key objects that allow for accurate measurements of both their mass and radius \\citep{2012A&A...539A..28G}. Equally important, they also allow for atmospheric measurements from transit transmission spectroscopy (e.g. \\citealt{2012ApJ...747...35B}) and/or occultation (secondary eclipse) emission spectroscopy \\citep{2012ApJ...751L..28D} without the challenging task of resolving the planet's light from the one of its host star. Such direct imaging is only possible to date for widely separated, massive young planets \\citep{2010Natur.468.1080M,2013A&A...559L..12A,2013ApJ...774...11K}. Transiting super-Earths around bright nearby stars are therefore essential to shed light on their true nature and origin.  HD\\,97658 is the second brightest ($V=7.7$, $K=5.7$) host star found to be transited by a super-Earth (the brightest being 55 Cnc; \\citealt{2011ApJ...737L..18W,2011A&A...533A.114D}). HD\\,97658\\,b was detected from radial velocity (RV) measurements with Keck-HIRES spectroscopy, with a minimum mass $M_{\\rm P} \\sin i = $ 8.2 $\\pm$ 1.2 $M_{\\oplus}$ and an orbital period $P=9.494 \\pm 0.005$ days \\citep{2011ApJ...730...10H}. The detection of transits from ground-based observations was announced \\citep{2011arXiv1109.2549H}, but follow-up space-based \\textit{MOST} photometry showed no transits at the claimed ephemeris \\citep{2012ApJ...759L..41D}. Further \\textit{MOST} observations finally discovered transit events with an ephemeris consistent with that predicted from RV measurements \\citep{2013ApJ...772L...2D}. From a global Markov Chain Monte-Carlo (MCMC) analysis simultaneously modeling HIRES RVs and \\textit{MOST} photometry, they derived a 2.34$^{+0.18}_{-0.15}\\,R_{\\oplus}$ and 7.86 $\\pm$ 0.73 $M_{\\oplus}$ planet, confirming its super-Earth nature with an average density of 3.44$^{+0.91}_{-0.82}$ g cm$^{-3}$.  The announcement of transit detections by \\citet{2013ApJ...772L...2D} motivated us to include HD\\,97658\\,b in our \\textit{Spitzer} program to search for and observe transits of low-mass RV-detected planets \\citep{2011IAUS..276..167G}. HD\\,97658\\,b was observed in August 2013 with the \\textit{Spitzer} Infra-Red Array Camera (IRAC, \\citealt{2004ApJS..154...10F}) at 4.5 $\\mu$m. \\textit{Spitzer} already proved to be the tool of choice to perform exquisite precise and continuous photometry for the detection and characterization of the ultra-shallow transits of super-Earths, such as HD\\,40307\\,b (for which transits are firmly discarded; \\citealt{2010A&A...518A..25G}) and 55 Cnc e (\\citealt{2011A&A...533A.114D}). We present here the confirmation of the transiting nature of HD\\,97658\\,b from \\textit{Spitzer} observations. We also present global Bayesian analyses of Keck-HIRES RV data, and \\textit{MOST} and \\textit{Spitzer} photometry, in order to derive as accurately as possible the parameters of HD\\,97658\\,b.  Ultimately, the precision for planetary parameters is related to the precision achieved for the stellar parameters, since planetary masses and radii as determined by transit and RV measurements cannot be determined independently from the properties of their host stars. Atmospheric stellar parameters such as effective temperature $T_{\\rm eff}$ and metallicity are determined from spectroscopy. The stellar radius, at least for nearby stars, can be derived from the luminosity (from parallax measurements) and the spectroscopic effective temperature using the classical equation of stellar physics $L_*=4\\pi R_*^2 \\sigma T_{\\rm eff}^4$. Obtaining the stellar mass is a more delicate process that can be achieved by a stellar evolution code using $T_{\\rm eff}$, metallicity and luminosity as inputs. The stellar age is usually poorly constrained in this process \\citep{2010ARA&A..48..581S}. A powerful method to provide accurate and precise stellar radius, mass and age is asteroseismology, by modeling the oscillation spectrum of the star. The CoRoT \\citep{2006ESASP1306...33B} and \\textit{Kepler} \\citep{2010ApJ...713L..79K} space missions were based on this complementarity, while the future TESS \\citep{2010AAS...21545006R} and PLATO \\citep{2013Rauer} space missions will also include significant asteroseismic programs. We also investigate here how future asteroseismic observations of HD\\,97658 can help to further improve the knowledge of the star and planet parameters.  This paper is organized as follows. The \\textit{Spitzer} observations and their reduction are presented in Sect.~\\ref{data}, while host star modeling is presented in Sect.~\\ref{host}. Section~\\ref{Spitzer} is devoted to the analysis of the \\textit{Spitzer} data, while Sect.~\\ref{Analysis} presents global Bayesian analyses of RV, \\textit{MOST} and \\textit{Spitzer} data. We discuss the results and provide avenues to further improve the characterization of HD\\,97658\\,b and its host star in Sect.~\\ref{discuss}. Conclusion and prospects are given in Sect.~\\ref{cc}. ",
        "conclusions": "\\label{cc} HD\\,97658\\,b is a key transiting super-Earth. The current data from \\textit{Spitzer} and \\textit{MOST} photometry, and also Keck-Hires RVs, analyzed with our MCMC code, indicate a planet with an intermediate average density ($\\rho_P=3.90^{+0.70}_{-0.61}$ g cm$^{-3}$). Investigating the possible internal compositions of HD\\,97658\\,b, our results indicate a large rocky component by at least 60\\% by mass, an amount of 0-40\\% of water+ices, and very little H-He components, at most 2\\% by mass. If interpreted as constraints on the atmospheric composition, the transit depths in both the \\textit{Spitzer} and \\textit{MOST} bandpasses together favor either a H$_2$-rich atmosphere with hazes or a cloud-free atmosphere with a low metallicity. Six more transit observations are planned in the course of 2014 with the \\textit{Spitzer} IRAC camera, this time with 3.6 $\\mu$m channel (PI: D. Dragomir).  HD\\,97658\\,b will be a target of the coming space missions TESS and CHEOPS, in particular to accurately measure the planet radius at visible wavelengths. By the time TESS and CHEOPS are launched (circa 2017), the GAIA mission\\footnote{http://sci.esa.int/gaia} (launched in December 2013) is hoped to have provided extremely accurate parallax measurements, within $\\sim$ 10 $\\mu$as for the bright nearby stars such as HD\\,97658. This will improve the knowledge on the distance of the star by 2 orders of magnitude compared to the current \\textit{Hipparcos} parallax. From there, the stellar radius will be known with an accuracy of 1\\% (assuming a $\\pm$ 50 K uncertainty on $T_{\\rm eff}$).  Knowing the stellar radius to 1\\% is also what we can expect for present and near-future long baseline interferometers \\citep{2013ApJ...771...40B}. This is the ultimate precision we can achieve on the planet radius, providing that the number of observed transits is sufficient to measure the transit depth $dF \\propto (R_P/R_*)^2$ with an accuracy much below 1\\%. ESA has now officially confirmed the selection of PLATO as the next M3 mission, to be launched around 2024. If HD\\,97658 falls in the observed fields, PLATO would not only measure the transit depth with unprecedented accuracy, but would also be able to detect oscillation frequencies of the host star that provide through asteroseismology accurate stellar mass and, very importantly, age.  Orbiting a bright K1-type star ($V=7.7$, $K=5.7$), HD\\,97658\\,b is a target of choice for atmospheric characterization. HST time is allocated to take in the coming months transmission spectroscopy of HD\\,97658\\,b with the WFC3 camera in IR (PI: H. Knutson). The aim is to distinguish a large scale-height, H-dominated atmosphere, from a compact, water steam atmosphere, which will confirm/refute the constraints on the atmospheric composition derived in this work. By comparing these measurements to the few targets for which such transmission spectroscopy is currently possible (55 Cnc e and GJ 1214\\,b), this will provide the first measure of the diversity of super-Earth atmospheres prior the era of the European Extremely Large Telescope (E-ELT) and the James Webb Space Telescope (JWST)."
    },
    "1402/1402.5899_arXiv.txt": {
        "abstract": "This paper quantitatively evaluates structural failure of biaxial cohesionless ellipsoids that have a two-density-layer distribution. The internal density layer is modeled as a sphere, while the external density layer is the rest of the part. The density is supposed to be constant in each layer. The present study derives averaged stresses over the whole volume of these bodies and uses limit analysis to determine their global failure. The upper bound condition of global failure is considered in terms of the size of the internal layer and the aspect ratio of the shape. The result shows that the two-density-layer causes the body to have different strength against structural failure. ",
        "introduction": "Theoretical studies about plastic analysis, using perfect ellipsoids, have been recently successful in explaining deformation processes of small bodies. \\cite{Holsapple2001} firstly used limit analysis to calculate the limit spins of cohesionless rubble pile ellipsoids and showed that any elements in a perfect ellipsoid reach the limit state at the same time. \\cite{Holsapple2004} reported that the limit spin given by volume averaged stresses is identical to the analysis given by \\cite{Holsapple2001}. Using the limit analysis approach, \\cite{Holsapple2006} investigated the tidal disruption condition of cohesionless ellipsoids. \\cite{Sharma2009}, on the other hand, introduced an averaging form over the whole volume of dynamical deformation to discuss structural failure of a rubble pile ellipsoid due to a tidal effect of a massive body. \\cite{Sharma2010} extended the theory by \\cite{Sharma2009} to binary-ellipsoid systems.  Regardless of their elegant mathematical formulations, these researches simplified their discussion by ignoring the effect of density distribution. It is interesting to understand how the density distribution has an effect on structural stability. In asteroid environments, the density may be distributed axisymmetrically during its accretion process. The present study considers the effect of this axisymmetric density distribution on the structural stability of asteroids. To investigate this effect, a uniformly rotating asteroid is modeled as a biaxial ellipsoid composed of an internal sphere and an external shell. The density is assumed to be constant in each layer. The technique used here is based on the limit analysis technique by \\cite{Holsapple2004} who obtained the limit spin by using the total volume stress. The current paper is organized as follows. First, the two-layer model is established. Second, the limit analysis technique is applied to the two-layer model. Last, the upper bound condition of structural failure of the present model is compared with that of the uniform density case. ",
        "conclusions": "This paper investigated the effect of the two-layer density distribution on structural failure of a uniformly rotating ellipsoid. The prime result shows that the two-layer density distribution causes different failure conditions from a uniformly rotating ellipsoid. The larger (smaller) size and higher (lower) density of the internal layer allow the bodies to become stronger (weaker) against structural failure. On the other hand, the critical spins with regards to the aspect ratio behave differently. If there is a high density core, the critical spins in the tension regime can increase, which indicates that the body becomes stronger against tension failure. However, in the compression regime, the critical spins for the density distribution case is  also larger than those for the uniform density case if $\\beta$ decreases. This implies that the body becomes weaker against compression failure. It can be explained that for such elongated bodies, since a dense core causes stronger gravitational compression and centrifugal forces do not effectively support the bodies, the bodies become sensitive to structural failure."
    },
    "1402/1402.0228_arXiv.txt": {
        "abstract": "We investigate radio emission efficiency $\\xi$ of pulsars and report a near linear inverse correlation between $\\xi$ and the spindown power $\\dot E$, as well as a near linear correlation between $\\xi$ and pulsar age $\\tau$. This is a consequence of very weak, if any, dependences of radio luminosity $L$ on pulsar period $P$ and period derivative $\\dot{P}$, in contrast to X-ray or $\\gamma$-ray emission luminosities. The analysis of radio fluxes suggests that these correlations are not due to a selection effect, but are intrinsic to the pulsar radio emission physics. We have found that, although with a large variance,  the radio luminosity of pulsars is $\\left<L\\right>\\approx 10^{29} \\,{\\rm erg/s}$, regardless of the position in the $P-\\dot P$ diagram. Within such a picture, a model-independent statement can be made that the death line of radio pulsars corresponds to an upper limit in the efficiency of radio emission. If we introduce the maximum value for a radio efficiency into Monte Carlo-based population syntheses we can reproduce the observed sample using the random luminosity model. The Kolmogorov-Smirnov test on a synthetic flux distribution shows high probability of reproducing the observed distribution. Our results suggests that the plasma responsible for generating radio emission is produced under similar conditions regardless of pulsar age, dipolar magnetic field strength, and spin-down rate. The magnetic fields near the pulsar surface are likely dominated by crust-anchored magnetic anomalies, which do not significantly differ among pulsars, leading to similar conditions for generating electron-positron pairs necessary to power radio emission. ",
        "introduction": "Despite the fact that pulsars were discovered almost half a century ago, the emission mechanism of the pulsed radio emission remains unresolved. The challenge lies in the difficulty to identify the correct ``coherent'' mechanism \\citep{2006_Melrose} that can power radio emission with similar emission characteristics over about four orders of magnitude in rotation period $P$ and up to eleven orders of magnitude in the rate of increase of pulsar periods $\\dot{P} = {\\rm d} P / {\\rm d} t$. In this paper we perform a systematic analysis of radio emission efficiency for the current sample of pulsars\\footnote{If not stated otherwise, data presented in this paper are taken from ATNF Pulsar Catalogue {\\it http://www.atnf.csiro.au/research/pulsar/psrcat} \\citep{2005_Manchester}}, in an effort to study the global radio emission properties and to constrain radio emission mechanism. In Section \\ref{sec:calc} we present the method to calculate radio luminosity $L$ and radio emission efficiency $\\xi$ of pulsars. The dependences of both $L$ and $\\xi$ on pulsar parameters are presented in Section \\ref{sec:res}. The results are summarized in Section \\ref{sec:sum} along with a discussion on the physical implications of the found correlations. \\newpage ",
        "conclusions": "\\label{sec:sum} Using a large sample of pulsars from the ATNF catalog, we draw the following conclusions in this paper: 1. Radio pulsar luminosity has a very weak, if any, dependence on pulsar spin-down luminosity $\\dot E$ over near 8 orders of magnitude in $\\dot E$. There might be a slight variation of the exponent of the $L-\\dot E$ correlation with frequency. 2. There is a near linear inverse correlation between radio efficiency $\\xi$ and $\\dot E$ for both normal and binary/millisecond pulsars. 3. The radio efficiency $\\xi$ is roughly linearly correlated with pulsar age $\\tau$ for normal pulsars. 4. The two reported correlations are not due to an observational selection effect. 5. Since radio luminosity does not depend on $P$ and $\\dot{P}$ (or the dependence is very weak) the proposed mechanism of radio emission must explain the significant increase of radio efficiency for low-$\\dot{E}$ pulsars.  A weak $L-\\dot E$ correlation was hinted by previous studies \\citep[e.g.][]{1993_Lorimer,2010_Ridley}. \\cite{2006_Malov} found a weak positive dependence $L \\propto \\dot{E}^{0.29}$ with a smaller sample (338 objects), which was selected by requiring the pulsars to have well-determined spectra, and/or well-determined distance. We adopt a much larger sample, which has the following advantages. First, a larger sample size can insure more reliable correlations. Second, many uncertainties, e.g. unknown moment of inertia of neutron stars, the influence of scintillation on pulsar flux density, or how well the one-dimensional line-of-sight cut through an emission beam could represent the entire beam, can be effectively canceled out. With our enlarged sample, the index of the correlation, if any, is much flatter than the ones found in the previous studies \\cite[see e.g.][]{2002_Arzoumanian, 2006_Malov}.  As discussed in \\S2.1, by using the simplified formula (\\ref{eq:lum}), we have neglected some complicated pulsar geometry factors, such as $\\alpha$, $\\beta$ (which defines the duty cycle $\\delta$ along with $P$, and $\\rho$). For individual pulsars, the derived $L$ should have a large error. Including the entire sample would cancel out most of these geometric factors, so that our discovered $\\xi-\\dot E$ and $\\xi-\\tau$ correlations would be intrinsic. One possible factor that is not fully canceled is the $P$-dependence. This is because both $\\rho$ and $\\delta$ depend on $P$. However, the explicit exponent of $P$-dependence is unknown. If we assume that the mechanism of radio emission is the same for both normal and millisecond pulsars, we find that by adding a dependence of $\\sim P^{-0.5}$ to the integrated luminosity, the $\\xi-\\dot E$ correlations for both normal and binary/millisecond pulsars become consistent with each other. Adding this period dependence results in the following luminosity relationships: $L_{400}\\propto \\dot{E}^{0.27}$, $L_{1400}\\propto \\dot{E}^{0.19}$, $L_{2000}\\propto \\dot{E}^0$. The found correlations are closer but still shallower than the one found by \\cite{2006_Malov}. With such a correction, the $\\xi-\\dot E$ correlation is slightly shallower, i.e. $\\xi \\propto \\dot E^{-0.81}$, but is still very significant. It has an even lesser impact on the $\\xi-\\tau$ correlation, i.e. $\\xi \\propto \\tau^{1.01}$ after correction. In a more general form, we can write that  $\\delta^{-1} \\sin^2{(\\rho/2)} \\propto P^p \\dot{P}^q $. Since in our analysis we are not using geometry information we cannot unambiguously define both $p$ and $q$. However, assuming similar radio luminosity of  normal and millisecond pulsars we can write that $ \\delta^{-1} \\sin^2{(\\rho/2)} \\propto P^{p} \\dot{P}^{-0.16(2p+1)}$. Note that even introducing some model-dependent value of $p$ does not change the general picture presented in this paper.  It would seem natural to assume that the spin-down parameters affect the radio luminosity (and hence the radio efficiency) of pulsars. Therefore, many authors \\citep[e.g.][]{1969_Ostriker,1975_Lyne, 1986_Stollman, 2006_Malov, 2006_Faucher} tried to define the luminosity law, however, the results vary greatly (depending on pulsar samples and used methods) and were inconclusive \\cite[see e.g.][]{2013_Bagchi}. The spin-down parameters determine the magnetic field strength at the light cylinder \\mbox{$B_{\\rm LC} \\propto P^{-2.5} \\dot{P}^{0.5}$}. Then, assuming dipolar configuration of magnetic field we can estimate the field strength at the stellar surface as \\mbox{$B_{\\rm d} \\propto (P \\dot{P})^{0.5}$} and the vacuum potential drop as \\mbox{$V\\propto B_{\\rm d} P^{-1} \\propto P^{-0.5} \\dot{P}^{0.5}$} \\citep{1975_Ruderman}. The main parameters that affects properties of pulsar radio emission are density and distribution of the electron-positron pair plasma. To estimate the plasma properties we need to use some specific model of acceleration (e.g. the vacuum gap, the space-charged limited flow, the partially screened gap models) and also to take into account various emission processes (e.g. curvature radiation, inverse Compton scattering). It was shown \\citep[see e.g.][]{2000_Zhang_a,2001_Hibschman, 2001_Hibschman_b} that the plasma density and distribution highly depend on some factors which cannot be estimated by the spin-down parameters. Thus, we cannot specify the properties of plasma responsible for radio emission using only $P$ and $\\dot{P}$. However, assuming dipolar configuration of pulsar magnetic field, some theoretical predictions suggest that the older the pulsar, the smaller the final pair multiplicity \\citep{2001_Hibschman, 2001_Hibschman_b}. It would suggest that some kind of dependence of the radio luminosity on the pulsar age could also be found. As we have shown it is not the case for the observed sample of radio pulsars (see Figure \\ref{fig:p_pdot_age.luminosity}).  In order to reproduce the observed sample of pulsars FK06 argued that in the absence of torque decay pulsar radio luminosity has to depend on $P$ and $\\dot{P}$ as \\mbox{$L\\propto P^{-1.5}\\dot{P}^{0.5}$} (or \\mbox{$L\\propto P^{-1.39}\\dot{P}^{0.48}$} with exponents optimized by \\citet{2013_Bates}) . However, the power-law luminosity model results in relatively low P-value of the K-S test for the flux distribution. We have found that replacing the modeled death line \\citep{1992_Bhattacharya} by the radio efficiency limit allows to reproduce the observed sample even with the luminosity model previously excluded by FK06, namely the random luminosity model. Introducing the radio efficiency limit into MC simulations results in much higher probabilities from the K-S test, allows to avoid a pileup of pulsars near the theoretical death line and, furthermore, explains the existence of a few pulsars observed in the so-called graveyard region. We argue that the facts that the observed luminosity does not depend on both $P$ and $\\dot{P}$, and that the random luminosity model of the intrinsic luminosity allows to reproduce the observed sample is a strong indication in favor of the random model.  It is very difficult to explain the spread of up to four orders of magnitude for both $L$ and $\\xi$ for a given spin-down luminosity only by statistical uncertainties. As we have mentioned above there must be some parameter/parameters other than $P$, $\\dot P$ that influence the radio efficiency of pulsars. The clear $\\xi-\\dot E$ linear inverse correlation and $\\xi - \\tau$ linear correlation for radio emission call for the following physical picture: The plasma responsible for generating pulsar radio emission must be produced under similar conditions regardless of pulsar age, dipolar magnetic field strength, and spin down rate. One possibility is that the pulsar polar cap region may have dominant crust-anchored magnetic anomalies, so that the near surface magnetic field configuration significantly deviates from the dipolar geometry \\citep{2002_Gil}. X-ray observations of old pulsars indeed revealed hot spots that are significantly smaller than the conventional polar cap \\citep{2005_Zhang,2008_Gil,2013_Szary}, which is consistent with having strong non-dipolar magnetic field at the surface due to crust-anchored local anomalies.%  The crust-anchored anomalies can be characterized by a parameter $b=B_{\\rm s}/B_{\\rm d}=A_{\\rm dp}/A_{\\rm bb}$, which describes the ratio of the actual value of the magnetic field at the surface $B_{\\rm s}$ to the dipolar component of the magnetic field $B_{\\rm d}$ \\cite[see e.g.][]{2002_Gil_b}. Here \\mbox{$A_{\\rm dp}\\approx 6.2 \\times 10^4 P^{-1} \\, {\\rm m^2}$} is the conventional polar cap area (i.e. calculated assuming pure dipolar configuration of the magnetic field), and $A_{\\rm bb}$ is the actual hot spot area (i.e. derived from X-ray observations). The sample of pulsars with an X-ray hot spot for which $b$ can be estimated is small. Nevertheless, with this small sample \\citep{2013_Szary}, we can clearly see a correlation between pulsar age and $b$ (see Figure \\ref{fig:bparameter}).  \\begin{figure}[h!] \\begin{center} \\includegraphics{f8} \\end{center} \\caption{Dependence of the $b$ parameter on pulsar age $\\tau$ (see text for more details). Dashed line correspond to the linear fit for all pulsars with detected radio emission. Note that for \\mbox{PSR J0108-1431} we used the single blackbody fit performed by \\cite{2009_Pavlov}.} \\label{fig:bparameter} \\end{figure}  The found correlation indicates that when a pulsar becomes older, its surface magnetic field becomes more dominated by the crust-anchored magnetic anomalies generated e.g. by the Hall drift \\cite[see][]{2013_Geppert, 2013_Vigano}. Such a crust-anchored anomaly does not depend on $\\dot E$ and age, which allows pulsars to produce enough electron-positron pairs at an old age to power radio emission.   The fact that pulsars near the graveyard tend to have a very high $\\xi$ also has a profound implications in understanding radio pulsar death. Traditionally, pulsar death line \\citep{1975_Ruderman,1993_Chen} was defined by the condition of production of electron-positron pairs, which depends on many uncertainties, including the near-surface $\\gamma$-ray emission mechanism \\citep{2000_Zhang} and the magnetic field configurations \\citep{2001_Gil}. The results presented in this paper offer a simpler interpretation to radio pulsar death. It is possible that the required physical condition to power radio emission is similar among pulsars in all ages, which requires a certain minimum spin-down power. When pulsars spin down slightly below this threshold, the radio emission mechanism simply cannot operate. This explains why high $\\xi$ pulsars are located near the death line. A similar conclusion was drawn by \\cite{2002_Arzoumanian} based on modelling ``of the birth properties and rotational, kinematic, and luminosity evolution of a Monte Carlo population of neutron stars'', even though the luminosity and spin-down rate scaling presented in their work, $L \\propto \\dot{E}^{0.5}$, does not reflect the actual $L-\\dot{E}$ relationship (see Figure \\ref{fig:one}a)."
    },
    "1402/1402.0534_arXiv.txt": {
        "abstract": "We study conditions under which carbon clusters of different sizes form and stabilize. {We describe an approach to equilibrium by simulating tenuous carbon gas dynamics to long times.} First, we use reactive molecular dynamics simulations to describe the nucleation of long chains, large clusters, and complex cage structures in carbon and hydrogen rich interstellar gas phases. We study how temperature, particle density, presence of hydrogen, and carbon inflow affect the nucleation of molecular moieties with different characteristics, in accordance with astrophysical conditions. We extend the simulations to densities which are orders of magnitude lower than current laboratory densities, to temperatures relevant to circumstellar environments of planetary nebulae, and to longtime (microsecond) formation timescales. We correlate cluster size distributions from dynamical simulations with thermodynamic equilibrium intuitions, where at low temperatures and gas densities, entropy plays a significant role. ",
        "introduction": "Carbon-based chemistry of the interstellar medium (ISM) is now an established discipline. Presence of large variety of organic molecules such as aldehydes \\cite{marcelino07}, alcohols \\cite{friberg88,marcelino07}, ethers \\cite{herbst06}, amides \\cite{hollis06}, sugar \\cite{hollis00}, acids \\cite{herbst06,gibb04}, amino acids \\cite{botta02,cronin93}, and nitriles \\cite{belloche08}, as well as many long-chain hydrocarbon compounds in the dense interstellar medium have been confirmed by infrared, radio, millimeter, and submillimeter spectroscopic measurement \\cite{Irvine1987,Ohishi1998, Winnewisser1999,Schutte1996,herbst09,hollis00}. Detection of ring species, including aromatic rings, raises the possibility that biologically significant molecules may exist in the ISM. Between 10\\%-30\\% of the interstellar carbon budget is thought to reside in polycyclic aromatic hydrocarbons (PAHs) \\cite{tielens}.  The presence of complex molecules, with typical spectral properties, can be used to precisely probe the environments in which they reside and provide a fairly faithful history of ISM molecular clouds. IR emission lines (in the range of 3-14 $\\mu$m) have been assigned to C-C and C-H vibration bands of a broad class of polycyclic aromatic hydrocarbons (PAHs) \\cite{salama08, Ehrenfreund, Duley, Salama, Ricks, herbst09, tielens}.  The simplest PAH, naphthalene (C$_{10}$H$_8$) has been found in carbonaceous chondrites \\cite{zare08}, suggesting extraterrestrial origin. Although no specific PAH molecule has yet been unambiguously identified in space, the infrared emission bands have been observed in nearly all astronomical objects, from highly ionized HII regions to ultraluminous infrared galaxies.  Moreover, the simplest PAH, naphthalene (C$_{10}$H$_8$) has been already found in carbonaceous chondrites \\cite{zare08}, suggesting its extraterrestrial origin.  The emission features at $7.0$, $8.5$, $17.4$ and $18.9$ $\\mu$m, present in the spectra of Tc1, NGC 7023 and NGC 2023 planetary and reflection nebulae, have been conclusively associated with the vibrational lines of regular carbon cages (C$_{60}$ and C$_{70}$) \\cite{cami10,sellgren10}.  Moreover, a large class of unsolved mysteries in astronomy, such as diffuse interstellar bands (DIB) \\cite{allamandolla85}, extended red emission (ERE) \\cite{witt84}, and the $2,175$ {\\AA} ``bump\" \\cite{2175} trace their origins to large carbon and/or carbonaceous molecules. Single and multi-layer graphene have strong plasmon excitation bumps in the vicinity of the $2,175$ {\\AA} feature \\cite{eberlein08}. However, more work in needed to establish potential connection between these spectral features in astrophysical conditions. These challenges can not be solved without a strong theoretical input. Development of a comprehensive and detailed theoretical models for the origin and evolution of complex organic molecules such as PAHs, fullerenes, and other ordered carbon structures may be crucially important for understanding of the role that prebiotic molecules played on early Earth.  Most theoretical studies of formation of larger aromatic carbon chains, rings, and cages in ISM have been done close to equilibrium and at relatively high particle densities. Naphthalene (C$_{10}$H$_8$) was shown to be potentially formed in ISM via a barrierless reaction of phenyl radical (C$_6$H$_5$) and vinylacetylene (C$_4$H$_4$) \\cite{parker12}. It was proposed that C$_{60}$ in ISM can form in a top-down manner in successive de-hydrogenation of PAHs into graphene \\cite{berne12}, and formation of defects (pentagons) necessary for the synthesis of C$_{60}$. There are a number of issues associated with this hypothesis of C$_{60}$ formation: 1) the circumstellar environment of Tc1 planetary nebula is distinctly hydrogen poor, 2) graphene is only synthesized in the laboratory and is not believed to exist in ISM, and 3) the proposed formation of carbon cages in ISM (via the 6 to 5 member rings) was observed in the laboratory only at far from equilibrium conditions.  To our knowledge, there are no available dynamical simulations of formation of molecules with aromatic rings at astrophysically-relevant conditions, which are not even properly known. Orderly growth of planar carbon structures (hexagon rings) from precursor C$_2$ molecules necessitates low-temperature conditions, where vibrational distortions of the structures during their growth are minimized \\cite{irle03}.  Under such conditions no significantly curved (cage-like) structures are supposed to form \\cite{irle03}. Fast inflow of carbon material is needed to form curved carbon clusters and fullerenes (pentagon rings, ring-stacks, ring fusion). In the laboratory, fullerenes are created with remarkable efficiency ($\\approx 40$ \\% in a condensing carbon arc) under less understood far from equilibrium conditions.  {In this work, we describe an approach to equilibrium by simulating tenuous carbon (and hydrogen) gas dynamics to exceeding long times. Different carbon molecular clusters and moieties which form represent steady state local equilibria and are important yardsticks for understanding future time dynamics and equilibration in astrophysical time.}   \\vspace{-2mm} ",
        "conclusions": "\\vspace{-2mm}  We study conditions under which large carbon structures form and stabilize in the gas phase in interstellar space by using reactive MD simulations. The formation mechanism of the large clusters and graphene type structures are investigated by a series of MD simulations. The influence of the temperature and particle density on the shape of structures are studied. It is found that at high temperature ($T=2,000-3,000$ K), fullerene type clusters are obtained, whereas at relatively low temperatures, graphene type sheets along with long chain molecules/structures are synthesized. Analytical expressions for the small chain molecules formation can be used to predict the time scale to formation of nanostructures in interstellar medium.  We determine the probability distributions of cluster sizes at different temperatures and initial C atom number and compare them with the gamma-distribution. The effect of hydrogen atom addition on the structure of carbon clusters is also studied. At high temperature, fullerene type clusters are found whereas short and branched chains molecules, terminated with hydrogen atoms, are observed at low temperature.  At higher density of hydrogen atoms, we find that small unsaturated carbon chains are formed. When the concentration of hydrogen are relative low graphene type sheet (terminated by hydrogen atoms) and small cluster are observed. Finally, we simulate the self-assembly of carbon structures, in non-equilibrium conditions of periodic carbon inflow, resembling the astrophysical conditions.  These studies help in better understanding the tenuous gaseous conditions in the ISM for synthesis of large carbon structures. We find that the system entropy plays a significant role at extreme conditions of low temperature and low density in hindering synthesis and stabilization of large carbon structures. In future studies, we intend to investigate how dynamical quantum tunneling and stellar irradiation can lead to new pathways of formation of carbon clusters. {UV irradiation and shock waves can ionize the gas and play significantly in molecular synthesis in the ISM. Quantum chemical calculations are necessary to address and understand the ionization process. When the quantum chemical calculations are limited by the system size, classical (semi-empirical) methods with improved force field can be used to understand the nucleation process.}"
    },
    "1402/1402.5916_arXiv.txt": {
        "abstract": " ",
        "introduction": "Over the next decade, large-scale structure (LSS) surveys will play an increasingly important role in the measurement of cosmological parameters and as a probe of initial conditions. In order to relate % late-time observables to the physics of the early universe, several sources of secondary non-linearities need to be understood (see fig.~\\ref{fig:schematic}). Reducing the theory error is essential if the full potential of future surveys is to be realized.\\footnote{The number of useful modes in galaxy surveys scales as the cube of the maximum wavenumber, $k_{\\rm max}$, at which the theoretical predictions can still be trusted.  Even a relatively modest gain in $k_{\\rm max}$ can therefore dramatically impact the scientific potential of galaxy surveys (but see~\\cite{Rimes:2005xs, Crocce:2005xz}).} Non-linearities in the gravitational evolution can be characterized by numerical N-body simulations~\\cite{Springel:2005mi} and, on sufficiently large scales, by perturbation theory~\\cite{Bernardeau:2001qr, Bernardeau:2013oda}. Less well understood are non-linearities in the {\\it biasing} between the clustering of galaxies and the underlying dark matter density.  \\vspace{0.2cm} \\begin{figure}[h!] \\centering \\includegraphics[scale =0.8]{Schematic.pdf} \\caption{Non-linearities in the gravitational evolution, in the biasing and in redshift space distortions~(RSD) complicate the relationship between the primordial initial conditions and large-scale structure observables.} \\label{fig:schematic} \\end{figure}   \\vskip 4pt The biasing problem is already visible in dark matter-only simulations, where it is reflected in the biasing of dark matter halos. On large scales, linear biasing has been shown to be a good approximation: \\beq \\delta_{h} = b_1 \\delta  \\ , \\eeq where $\\delta_h$ and $\\delta$ are the density contrasts of the halos and the dark matter, respectively, and the bias parameter $b_1$ is an unknown coefficient (to be fit to data).  However, linear biasing is known to fail on small scales where non-linearities becomes important. One common procedure for describing halos beyond the linear biasing model is {\\it local Eulerian biasing}~\\cite{Fry:1992vr} which assumes that the halo density contrast is a local function of the dark matter density, $\\delta_h(\\x,\\tau)={\\cal F}[\\delta(\\x, \\tau)]$. Formally, we might write this relation as a Taylor expansion \\beq \\delta_h(\\x, \\tau) = \\sum_{n=0}^\\infty  \\frac{b_n^{(0)}}{n!} \\, \\delta^n(\\x, \\tau)\\label{equ:Fry} \\ . \\eeq Local biasing is motivated both as a natural generalization of linear biasing and as a consequence of a number of semi-analytic models of halo formation. It is also often employed in data analysis~\\cite{Gaztanaga:1994us, Scoccimarro:2000sp, Feldman:2000vk, Verde:2001sf, Gaztanaga:2005an, McBride:2010zp}. However, the meaning of (\\ref{equ:Fry}) is far from clear, as we need to define $\\delta^n(\\x, \\tau)$ for $n>1$. In particular, non-linear quantities like $\\delta^n$ receive contributions from all scales and are therefore not necessarily small, even on large scales. A common procedure is to smooth $\\delta$ with a window function that removes power below some length scale~$\\Lambda^{-1}$.  However, this only masks the problem. The scale~$\\Lambda^{-1}$ is arbitrary and should not appear in physical quantities.\\footnote{Semi-analytic models only add to the confusion, as they often identify $\\Lambda^{-1}$ with the Lagrangian size of the halo, which is a physical scale.} In~\\cite{McDonald:2006mx}, McDonald showed how to reorganize the bias expansion in a way that makes this property manifest (see also~\\cite{McDonald:2009dh, Schmidt:2012ys}). In this paper, we revisit his idea of {\\it renormalized halo bias}.  \\vskip 4pt The key feature of renormalization is that, although $\\delta^n$ may be large (or depend on an unphysical smoothing scale $\\Lambda$), all large (or $\\Lambda$-dependent) contributions can be systematically removed by adding local counterterms % \\beq [\\delta^n](\\x, \\tau) \\equiv \\delta^n(\\x, \\tau) + \\sum_{\\widetilde {\\cal O}} Z^{(\\delta^n)}_{\\widetilde{\\cal O}} \\hskip 1pt\\widetilde{\\cal O}(\\x, \\tau)  \\, \\ll\\, 1\\ . \\eeq We will find that consistent renormalization requires additional fields $\\widetilde {\\cal O}$ beyond those appearing in the local Eulerian biasing model~(\\ref{equ:Fry}). For instance, at quadratic order, we have to add the tidal tensor~$(\\nabla_i \\nabla_j \\Phi_g)^2$~\\cite{McDonald:2009dh, Chan:2012jj, Baldauf:2012hs}, where $\\Phi_g$ is the gravitational potential, while, at cubic order, the velocity potential $\\Phi_v$ has to be introduced as an independent degree of freedom.  In addition, higher-derivative terms must be included to remove the sub-leading $\\Lambda$-dependence.  Finally, no stochastic parameters are needed to remove divergences in correlation functions at separated points, although one may choose to include them to remove delta-function localized divergences.  In general, we find that all the terms that are allowed by symmetry are required in order to define the renormalized fields~$[{\\cal O}] = \\{ [\\delta^n],[\\widetilde{\\cal O}]\\}$.\\footnote{In this paper, we will only consider the case of Gaussian initial conditions. Additional terms can arise for non-Gaussian initial conditions~\\cite{McDonald:2008sc}. We will discuss this elsewhere~\\cite{NG-progress}.}  Renormalized halo biasing is naturally an {\\it effective theory}, in which the halo density contrast is written as a double expansion in terms of powers of the fluctuations and their derivatives.  In order to be well-defined, this expansion % has to be organized in terms of the renormalized quantities: \\beq \\delta_h(\\x, \\tau) \\, =\\,  \\sum_{\\cal O} b_{{\\cal O}}^{(R)} \\, [{\\cal O}](\\x, \\tau)\\ ,\\label{equ:dh} \\eeq where $b_{{\\cal O}}^{(R)}$ are the {\\it renormalized bias parameters} which by construction do {\\it not} depend on the smoothing of the density field. It is to be expected that the biasing model (\\ref{equ:dh}) will contain all terms~${\\cal O}$ consistent with the symmetries of the dark matter equations of motion, as was emphasized by McDonald and Roy~\\cite{McDonald:2009dh}~(see also \\cite{Kehagias:2013rpa}).\\footnote{The form of~(\\ref{equ:dh}) was also argued to arise from the non-linear time evolution of local Lagrangian biasing~\\cite{Chan:2012jj}.} After renormalization, terms with higher number of fields and higher number of spatial derivatives are suppressed.  To describe the halo statistics to a finite precision then only requires a finite number of terms in (\\ref{equ:dh}). How many terms need to be retained depends on the power spectrum of the initial conditions in the same way as in the effective field theory of large-scale structure (EFT-of-LSS)~\\cite{Carrasco:2012cv, Pajer:2013jj, Carrasco:2013mua}. In principle, the different terms in the bias expansion can be distinguished by measuring correlation functions of the dark matter halos and fitting for the bias parameters $b_{{\\cal O}}^{(R)}$.  For the gravitational tidal tensor this has been demonstrated in \\cite{Baldauf:2012hs, Sheth:2012fc}. Here, we lay out the basic steps towards a systematic treatment of the effective theory of halo biasing.   \\vskip 6pt The outline of the paper is as follows: In Section~\\ref{sec:Renorm}, we give a pedagogical introduction to renormalization in structure formation. We present the renormalization of $\\delta^2$ as an explicit example. In Section~\\ref{sec:Renorm2}, we generalize this to all terms in the local Eulerian biasing model.  We classify the terms that have to be added to the biasing model to make it consistent with renormalization. In Section~\\ref{sec:HaloStatistics}, we show how these terms affect halo statistics at one-loop order. Our conclusions appear in Section~\\ref{sec:Conclusions}. Four appendices contain technical details: In Appendix~\\ref{sec:mixing}, we discuss how the renormalized quantities depend on the renormalization scale. In Appendix~\\ref{ap:HOO}, we present the renormalization of $\\delta^n$, for $n > 2$. In Appendix~\\ref{app:Galileon}, we prove a non-renormalization theorem for Galileon operators. In Appendix~\\ref{app:Bispectrum}, we give the details of the one-loop halo bispectrum.  \\vskip 6pt We will use the following notation and conventions: Conformal time is denoted by $\\tau$ and comoving coordinates by $\\x$. Overdots denote derivatives with respect to conformal time and $\\H \\equiv \\dot a/a$ is the comoving Hubble parameter. For momentum integrals, we use the shorthand $\\int_\\p \\equiv \\int \\d^3 p/(2\\pi)^3$. We define the normalized inner product of two vectors as $\\mu_{\\p,\\q}\\equiv \\p\\cdot\\q/pq$ and sometimes use $\\sigma^2_{\\p,\\q} \\equiv \\mu^2_{\\p,\\q} - 1$. A prime on a correlation function, $\\langle \\ldots \\rangle'$, indicates that the overall delta function is dropped.  We will use $P(q) \\equiv \\langle \\delta^{(1)}_\\q \\delta^{(1)}_{\\q'}\\rangle'$ for the linear power spectrum (which is often called $P_{11}$) and $P_{mm}(q) \\equiv \\langle \\delta_\\q \\delta_{\\q'}\\rangle'$ for the non-linear power spectrum.  Sometimes, we write $P_i \\equiv P(q_i)$. Renormalized operators are denoted by square brackets, $[{\\cal O}]$.  We use two definitions for the bias parameters of $\\delta^n$, namely $b_n$ and $b_{\\delta^n}$, which are related by $b_{\\delta^n}\\equiv b_n/n!$. When we present numerical results, the linear power spectrum is computed with CAMB~\\cite{Lewis:1999bs}, using a $\\Lambda$CDM cosmology with the standard cosmological parameters, $\\Omega_{\\Lambda} = 0.73$, $\\Omega_b h^2 =0.02$, $\\Omega_m h^2 = 0.12$, $h = 0.7$, and power law initial conditions, with $n_s=0.96$ and $A_s = 2.2 \\times 10^{-9}$. ",
        "conclusions": "\\label{sec:Conclusions}  In this paper, we have shown explicitly how renormalization {\\it forces} us to treat biasing models as {\\it effective theories}, i.e.~as a double expansion in terms of fluctuations in the dark matter density and velocity and derivatives thereof. Consistently removing the short-scale physics in the local Eulerian biasing model~\\cite{Fry:1992vr} doesn't just renormalize the bias parameters, but also generates {non-local} terms and {higher-derivative} contributions.  In order for the theory to become independent of the unphysical regulator of composite operators, all terms consistent with the symmetries have to be included in the biasing model. At lowest order, this means adding the gravitational tidal tensor, while at cubic order, the {velocity potential} appears as an independent degree of freedom.  In the process, we have clarified a few technical aspects of the renormalization procedure proposed by McDonald~\\cite{McDonald:2006mx}. We organized the renormalization of composite operators in a convenient diagrammatic representation and derived the building blocks for the one-loop renormalization of the halo power spectrum and bispectrum. We proved that Galileon operators aren't renormalized at leading order in derivatives. Finally, we showed explicitly how the definition of the renormalized theory depends on the renormalization scale and how the terms in the bias expansion mix as this scale is varied.  This scale-dependence of the renormalized halo bias is relevant for interpreting recent N-body results~\\cite{Chan:2012jx, Sheth:2012fc, Pollack:2013alj}.   \\vskip 4pt Our work motivates several future directions:  \\begin{itemize} \\item Most importantly, it remains to be quantified how many terms need to be kept in the effective theory to achieve a given target of precision in the predictions for the statistics of halos. The importance of including the gravitational tidal tensor has been established through N-body simulations in \\cite{Baldauf:2012hs,Chan:2012jj}.  However, a completely systematic exploration of the effective theory of halo biasing has not yet been performed. For example, the scale $\\Lambda_\\star$ which suppresses higher-derivative terms has not been measured and the errors that arise from truncating the effective theory have not been quantified.  \\item A key motivation for understanding non-linear biasing is primordial non-Gaussianity. One may hope that the non-linearities arising from the biasing are sufficiently distinct, so that the primordial signals can be extracted from the shape information of the correlation functions.  To analyze this self-consistently, we must consider the effects of these non-Gaussian contributions on the renormalization of the biasing model~\\cite{McDonald:2008sc, NG-progress}.  \\item We have focused on Eulerian biasing, but it has been suggested that biasing in Lagrangian space may have some advantages.  The approach taken here can likely be adapted to the Lagrangian EFT-of-LSS \\cite{Porto:2013qua}.  Which scheme is more useful may depend on the type of observable that is considered.  \\item  We have shown how renormalization forces biasing to be non-local in space, but the theory has remained local in time. Recently, the necessity of non-locality in time has been emphasized in the EFT-of-LSS~\\cite{Carrasco:2013sva, Carroll:2013oxa}. This deserves further consideration in the present context.  \\item Ultimately, galaxies are observed in redshift space. Redshift space distortions can add extra non-linear contributions to the observed galaxy correlation functions. This should be taken into account.  \\item  We have seen that a biasing model in terms of just the dark matter density  is not consistent and that the velocity potential has to be added as independent degree of freedom.   It would be interesting to explore the physical effects of the velocity potential more widely, e.g.~beyond perturbation theory. \\end{itemize}    \\subsubsection*{Acknowledgements}  We thank Tobias Baldauf, Hayden Lee, Enrico Pajer, Marcel Schmittfull, Leonardo Senatore and Marko Simonovi\\'c for helpful discussions. D.B.~and V.A.~gratefully acknowledge support from a Starting Grant of the European Research Council (ERC STG grant 279617).  The research of D.G.~is supported in part by the Stanford Institute for Theoretical Physics and by the U.S. Department of Energy contract to SLAC no.\\ DE-AC02-76SF00515. M.Z.~is supported in part by the NSF grants PHY-0855425, AST-0907969, PHY-1213563 and by the David and Lucile Packard Foundation.  \\newpage \\appendix"
    },
    "1402/1402.1829_arXiv.txt": {
        "abstract": "We apply the statefinder diagnostic to the holographic dark energy models, including the original holographic dark energy (HDE) model, the new holographic dark energy model, the new agegraphic dark energy (NADE) model, and the Ricci dark energy model. In the low-redshift region the holographic dark energy models are degenerate with each other and with the $\\Lambda$CDM model in the $H(z)$ and $q(z)$ evolutions. In particular, the HDE model is highly degenerate with the $\\Lambda$CDM model, and in the HDE model the cases with different parameter values are also in strong degeneracy. Since the observational data are mainly within the low-redshift region, it is very important to break this low-redshift degeneracy in the $H(z)$ and $q(z)$ diagnostics by using some quantities with higher order derivatives of the scale factor. It is shown that the statefinder diagnostic $r(z)$ is very useful in breaking the low-redshift degeneracies. By employing the statefinder diagnostic the holographic dark energy models can be differentiated efficiently in the low-redshift region. The degeneracy between the holographic dark energy models and the $\\Lambda$CDM model can also be broken by this method. Especially for the HDE model, all the previous strong degeneracies appearing in the $H(z)$ and $q(z)$ diagnostics are broken effectively. But for the NADE model, the degeneracy between the cases with different parameter values cannot be broken, even though the statefinder diagnostic is used. A direct comparison of the holographic dark energy models in the $r$--$s$ plane is also made, in which the separations between the models (including the $\\Lambda$CDM model) can be directly measured in the light of the current values $\\{r_0,~s_0\\}$ of the models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3701_arXiv.txt": {
        "abstract": "We use a parent sample of 118 gamma-ray burst (GRB) afterglows, with known redshift and host galaxy extinction, to separate afterglows with and without signatures of dominant reverse-shock emission and to determine which physical conditions lead to a prominent reverse-shock emission. We identify 10 GRBs with reverse shock signatures - GRBs 990123, 021004, 021211, 060908, 061126, 080319B, 081007, 090102, 090424 and 130427A. By modeling their optical afterglows with reverse and forward shock analytic light curves and using Monte Carlo simulations, we estimate the parameter space of the physical quantities describing the ejecta and circumburst medium. We find that physical properties cover a wide parameter space and do not seem to cluster around any preferential values. Comparing the rest-frame optical, X-ray and high-energy properties of the larger sample of non-RS-dominated GRBs, we show that the early-time ($<$ 1ks) optical spectral luminosity, X-ray afterglow luminosity and $\\gamma$-ray energy output of our reverse-shock dominated sample do not differ significantly from the general population at early times. However, the GRBs with dominant reverse shock emission have fainter than average optical forward-shock emission at late time ($>$ 10 ks). We find that GRBs with an identifiable reverse shock component show high magnetization parameter $R_{\\mathrm{B}} = \\varepsilon_{\\rm B,r}/\\varepsilon_{\\rm B,f} \\sim 2 - 10^4$. Our results are in agreement with the mildly magnetized baryonic jet model of GRBs. ",
        "introduction": "\\label{s1} The study of gamma-ray burst (GRB) afterglows started with their discovery in 1997 \\citep[][]{costa1997,paradijs1997}. Since then, afterglow observations have established the cosmological nature of GRBs \\citep[e.g.,][]{gomboc2012}, provided information on their stellar progenitors \\citep{hjorth2012} and prompted the study of GRB circumburst environment \\citep[e.g., ][]{petitjean2011}, their host galaxies \\citep[][]{berger2011, fynbo2012}, and intergalactic medium in the GRB line-of-sight \\citep[e.g., ][ and references therein]{vergani2009, chornock2013}.  Afterglow emission has also been considered as a powerful probe capable of revealing physical properties in gamma-ray burst ejecta as well as the medium through which the fireball propagates. According to the standard afterglow model \\citep{sari1995,meszaros1997}, a relativistically expanding fireball propagating through a medium surrounding a GRB progenitor drives a shock into the medium, known as a forward shock (FS). Heated electrons behind the shock emit synchrotron radiation, giving rise to FS afterglow emission \\citep{sari1998}. In addition to the forward shock, a reverse shock (RS), propagating back into the fireball, can be produced \\citep{sari1999b}.  The FS afterglow model has proven to describe late-time afterglow behavior well.  The environmental dependence of afterglow light curve evolution has been calculated analytically for two limiting cases: constant density interstellar medium (ISM) \\citep[e.g., ][]{sari1998} and stellar wind environment with a density profile $\\propto r^{-2}$ \\citep[e.g., ][]{chevalier2000}, where $r$ is distance from the progenitor. The distinct temporal and spectral behavior of the two environments can be used to determine the nature of the medium surrounding the progenitors \\citep{schulze2011} using predicted relations between  temporal and spectral afterglow slopes, i.e., closure relations \\citep[e.g., ][]{zhang2006}. However,  in order to constrain the values of the microphysical parameters in the ejecta (the ratio of the electron and magnetic energy density over the internal energy density in the shocked region, $\\varepsilon_{\\mathrm{e}}$ and $\\varepsilon_{\\mathrm{B}}$, electron energy distribution index $p$), an analysis of the FS emission alone is insufficient. Emission from the RS can be used to measure the values of microphysical parameters,  to derive a Lorentz factor of the ejecta, and to constrain the nature of the ejecta itself \\citep{zhang2003,kobayashi2003}.  The evolution of the RS emission has been mostly studied for two limiting cases \\citep{kobayashi2000}: $(i)$ the thick shell case, in which a RS becomes relativistic and starts to decelerate shell material, and $(ii)$ the thin shell case, in which a RS stays sub-relativistic and is too weak to decelerate the shell. Observational evidence suggest that the two limiting cases might not describe the real conditions very well \\citep[e.g.,][]{virgili2013}. Intermediate events (between the relativistic and sub-relativistic cases) should be handled by numerical studies \\citep{nakar2004,harrison2013}. All these studies assume a brief central engine activity and a constant Lorentz factor of the ejecta, resulting in a \\textit{short-lived} reverse shock emission. However, ejecta could be composed of shells of different Lorentz factors. In such scenario, slower shells continue to feed the blast wave, giving rise to a \\textit{long-lived} reverse shock emission. In this case, depending on the mycrophysics parameters, light curve can be completely dominated by a RS emission for the duration of the afterglow \\citep[][]{uhm2007,genet2007}. This work focuses exclusively on the short-lived RS emission.  RS emission is expected to be especially prominent at low frequencies (optical to radio) and can be recognized by its characteristic rising and decaying slopes. For a typical set of micropysical parameters and initial Lorentz factor of the shell, the optical band lies between the typical ($\\nu_{m,r}$) and cooling ($\\nu_{c,r}$) synchrotron frequencies of the RS. In a constant ISM medium, the emission is predicted to reach its peak at the fireball deceleration time and then decay with a characteristic power-law slope of $\\sim 2$ \\citep{kobayashi2000} for both thin- and thick-shell cases. The rising index should be very steep ($\\sim -5$) for a thin-shell or shallow ($\\sim -1/2$) in the thick-shell approximation. In the case of a wind medium, a shallow rise and a steeper decay slope of $\\sim 3$ are expected (for standard parameters) for both thick- and thin-shell scenarios \\citep{kobayashi2003b, zou2005}. Depending on the relative strength and peak times of a reverse and forward shock afterglow component, three different light curve configurations are expected to be observed \\citep[][]{zhang2003,gomboc2009}: \\begin{itemize} \\item Type I: light curve with prominent reverse and forward shock afterglow peaks, \\item Type II: light curve with characteristic flattening due to bright RS afterglow outshining the FS emission, \\item Type III:  light curve with simultaneous FS and RS peaks, where the former outshines the latter. \\end{itemize} RS emission arising from mildly magnetized outflows is predicted to be highly polarized \\citep{lyutikov2003, granot2003, rossi2004}. Polarimetric measurements of early-time afterglow emission can thus provide a complementary method of recognizing or confirming a RS emission component \\citep{mundell2007a,steele2009,mundell2013}.  Before 2005, only a few afterglows had been observed less than $\\sim 1$ hour after the GRB trigger, with RS components being identified in three of them (GRB\\,990123 - e.g., \\citealt{sari1999a}, \\citealt{kobayashi2000}, \\citealt{soderberg2002}; GRB\\,021004 - e.g. \\citealt{kobayashi2003}, \\citealt{kobayashi2003b}; GRB\\,021211 - e.g., \\citealt{fox2003}, \\citealt{wei2003}).  With the launch of NASA's \\textit{Swift} satellite \\citep{gehrels2004} and the advent of purpose-built rapid-response autonomous robotic telescopes (such as the Liverpool Telescope and Faulkes telescopes [LT/FT; \\citealt{steele2004}], Rapid Eye Mount [REM; \\cite{zerbi2001}], Robotic Optical Transient Search Experiment [ROTSE; \\citealt{akerlof2003}], etc.), the possibility of routinely observing the very early afterglow of large numbers of GRBs became a reality. RS optical flashes were expected to be ubiquitous in this new era of rapid follow-up. Surprisingly, the rate of detected RS components has been extremely low \\citep{roming2006,melandri2008,gomboc2009,oates2009}. The lack of detections has been attributed either to strongly magnetized outflows, which can suppress the RS emission \\citep{zhang2005}, the RS emission peaking at lower frequencies than the optical band (e.g. IR/mm) at early time \\citep{mundell2007b,melandri2010}, or prompt optical emission originating in an internal shock region outshining  any contemporaneous external RS emission \\citep{kopac2013}.  To better understand the nature of RS emission, we compile a sample of GRBs with optical afterglows which show RS signatures. We compare their rest-frame optical, X-ray and $\\gamma$-ray properties to a larger sample of GRBs with no apparent RS contribution in their optical afterglow. To investigate whether the physical conditions in the GRB ejecta of our RS afterglow sample show similar or different properties, we use a simple analytical model of reverse- and forward-shock emission and apply it to our RS sample using Monte Carlo simulations. We examine a parameter space that describes our afterglows well and discuss relations between various parameters.  Throughout the paper the convention $F_{\\nu,\\mathrm{t}} \\propto t^{-\\alpha}\\nu^{-\\beta}$ is adopted, where $\\beta$ and $\\alpha$ are spectral and temporal afterglow slopes, respectively. Standard cosmology with $H_{0} = 71~\\mathrm{km}~\\mathrm{s}^{-1}~\\mathrm{Mpc}^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ is assumed.  Times are given with respect to GRB trigger time. ",
        "conclusions": "In this work we present a detailed study of a sample of 10 GRB afterglows that show RS signatures in early optical light curves. The sample is composed of one Type I (in which both reverse and forward shock afterglow light curve peaks are observed) and  nine Type II light curves (in which the characteristic steep-to-shallow light curve evolution, caused by the dominant RS at early time and the later rise of FS emission, is observed), as classified in \\citet{zhang2003} and \\citet{gomboc2009}. The 10 afterglows represent only a fraction of a much larger sample, composed of 118 afterglows with measured redshift and host galaxy extinction, which we compiled in order to investigate the rest-frame properties of the former in relation to the larger sample.  We compare the rest-frame optical, X-ray and $\\gamma$-ray properties of the RS sample to a sample of afterglows without compelling RS signatures. Early-time RS emission is found to span over several orders in spectral luminosity, which is consistent with the general early-time spread in afterglows' brightness \\citep[e.g., ][]{kann2010}. On the other hand, we find that all but one afterglow from our sample are among the faintest at late times ($t_{\\mathrm{rest}} > 10$ ks). Since only 6 out of 10 RS afterglows were observed with the \\textit{Swift} XRT instrument, we cannot draw any firm conclusions about the X-ray properties of the sample. The high-energy properties (i.e., isotropic equivalent energy) of RS and non-RS GRBs in our full sample do not statistically differ.  Using a simple analytic model of a RS and FS afterglow we reproduce the observed optical light curves of our RS sample by using a MC simulation. Derived physical properties do not reveal any preferential values within the assumed parameter space. This is similar to the results obtained in previous analyses which concentrated on late time FS emission \\citep[e.g., ][]{panaitescu2001,panaitescu2002}. Failure to reproduce X-ray and radio observations, where available, points to the need to either change the basic assumptions of the model (e.g., thick- vs. thin-shell scenario, ISM vs. stellar wind circumburst medium) or introduce more complicated emission components beyond the simple standard theory.  According to \\citet{zhang2005}, a strong RS emission is produced when the GRB outflow is baryonic, i.e., only mildly magnetized. Furthermore, in order to produce a RS afterglow that can outshine the FS emission (Type II light curve), a magnetization parameter of $R_{\\rm B} > 1$ is required. We find that our RS sample afterglows have preferentially both low $\\varepsilon_{\\rm B,r}$ as well as high $R_{\\rm B}$ values. Consequently,  the presence of strong RS emission (compared to  FS emission) requires $\\varepsilon_{\\rm B,f} < \\varepsilon_{\\rm B,r}$. In the standard FS afterglow model, the peak in the spectral domain $F_{\\rm max,f}$ is proportional to $\\varepsilon_{\\rm B,f}^{1/2}$ \\citep{sari1998}. Thus, a low value of $\\varepsilon_{\\rm B,f}$ is expected to produce fainter FS emission, which is what we find in our RS sample at optical wavelengths. In addition, the time of the FS peak is proportional to $t_{\\rm p} \\propto \\varepsilon_{\\rm B,f}^{1/3}$ \\citep{sari1998}. This could explain the lack of Type I light curves, since the FS peak for low $R_{\\mathrm{B}}$ ratio is likely to occur when the RS afterglow component is still very bright. Due to different models used in the literature as well as our mostly unconstrained values of $\\varepsilon_{\\rm B,f}$, we cannot test whether the derived $\\varepsilon_{\\rm B,f}$ in afterglows with prominent RS components is generally lower than in non-RS events. The interpretation of faint late-time optical afterglows in our RS sample may be revised if there is an intrinsic correlation between $\\varepsilon_{\\rm B,f}$ and other parameters that define the afterglow emission (e.g., \\citealt{santana2013} recently found a hint of correlation between parameters $\\varepsilon_{\\rm B,f}$ and $E_{\\rm K}$).  Fifteen years after the discovery of GRB\\,990123, it is clear that larger samples of confirmed RS components are vital to understand the nature of RS emission and to determine the origin of RS suppression. In addition to the standard techniques involving light curve and spectral analysis, unambiguous identification of RS components may become increasingly possible via the detection of early time optical polarization \\citep{mundell2007a,steele2009,mundell2013} with simultaneous multicolor light curves using new polarimeters on robotic telescopes, such as RINGO3, mounted on the Liverpool Telescope \\citep{arnold2012}."
    },
    "1402/1402.6062_arXiv.txt": {
        "abstract": "A compensator (null-corrector) for testing aspheric optical surfaces is proposed, which enables {\\it i)}~independent verification of optical elements and assembling of the compensator itself, and {\\it ii)}~ascertaining the compensator position in a control layout for a specified aspheric surface. The compensator consists of three spherical lenses made of the same glass. In this paper, the scope of the compensator expanded to a surface speed $\\sim f/2.3$; a conceptual example for a nominal primary of Hubble Space Telescope is given. The autocollimating design allows significant reducing difficulties associated with practical use of lens compensators. ",
        "introduction": "As is known, it is easy to control a spherical surface during its making by examining the image of a point light source disposed at the center of curvature. If the sphere is perfect, we should see the diffraction Airy pattern; deviations from the desired surface shape can be evaluated either qualitatively, by using the Foucault knife-edge test, or quantitatively, by imposing a reference wavefront on the studied wave and analyzing the resulting interferogram.  On the contrary, testing of aspheric surfaces confronts opticians with a much more complicated problem. The point is that the spherical wave, being reflected from an aspheric surface, skews its original shape and does not form any reasonable image. Let us suppose, for example, that we have to control a surface of revolution of a conic section of diameter $D$ with the paraxial curvature radius $R_0$ and squared eccentricity $\\varepsilon^2$. Unlike the sphere, a bundle of normals to this surface do not converge in a single point; the distance $N(y)$ between points of convergence of the paraxial normals and those to the $y$-zone -- {\\it an aberration of normals} -- is defined by a simple expression: \\begin{equation} N(y) = \\varepsilon^2 s(y), \\end{equation} where $s(y)$ is sagitta for the $y$-zone (we assume, as usual, that the $z$-axis is directed along the system's axis of symmetry). E.g., in the case of the nominal primary mirror of the Hubble Space Telescope (HST) we have $D = 2400$~mm, $R_0 = 11040$~mm, $\\varepsilon^2 = 1.002299$, so the marginal sagitta $s(D/2) \\simeq 65.22$~mm, and the marginal normals aberration $N(D/2) \\simeq 65.37$~mm. If we put a point light source in the paraxial curvature center, then the diameter of the reflected beam in the source vicinity exceeds $28$~mm, and the image-based control of the mirror is simply out of the question.  Under these conditions, aspheric surfaces are usually controlled indirectly. Namely, a studied aspheric is included as a component into some more extent optical system in such a way that the wavefront emerging the whole system became spherical. The easiest way is to compensate the divergence of normals to aspheric surface. The corresponding methods, dating back to Maksutov (1924, 1932; see Maksutov 1984, p. 237) and \\citet{b1}, have now become basic at control of astronomical optics. The theory of compensators is discussed in detail in monographs of Wilson (1999, Sec.~1.3.4) and Geary (2002, Ch.~35), as well as in numerous papers mentioned there.  Clearly, tight tolerances are inherent to compensators, but not this feature lies in the heart of the problem; quite similar tolerances are inherent to some other optical systems. Having high enough optical power, a compensator must bring into a wavefront huge spherical aberration, moreover, of strictly specified value. {\\em Meanwhile, the conventional compensator cannot be verified and properly positioned without the use of extraneous optics.} Thus, we must distinguish errors of the tested surface from those of the compensator and auxiliary devices. For this reason, leading optical workshops use a few compensators of different type, and only at coincidence of the analysis results the surface shape can be considered as ascertained.  Let us note in this regard that a known fault at figuring the HST primary mirror associated with wrong alignment of two-mirror compensator was essentially due to neglecting results obtained with the other, two-lens compensator of \\citet{b7}. According to Wilson (1999, p.~85), equally significant errors were made at manufacturing of mirrors for large ground-based telescopes, but they have not caused wide public response.  The need for reliable inferences forces us to pay special attention to possibilities of {\\em independent} checking a compensator itself and its position in the control scheme. A few layouts partially satisfying to these requirements are described in the literature \\citep{b8,b5,b13}. This paper represents an {\\em autocollimating compensator}, a device that meets requirements mentioned above. As a consequence, one can definitely refer just to the aspheric surface all defects of the wavefront remaining visible after performance of specified procedures.  A particular scheme of the autocollimating compensator has been proposed several years ago \\citep{b10} in connection with the prospective renovation of the G.A.~Shain 2.6-m telescope in the Crimean Astrophysical Observatory (CrAO, Ukraine). Then the compensator was manufactured in  CrAO optical workshop under the direction of N.V. Steshenko. A more detailed description of that layout has been recently given by \\citet{b12}. It was noted in the latter paper that the compensator to a faster aspheric mirror than the $f/3.85$ G.A.~Shain primary is of considerable interest. Just this case is discussed below on an example of the $f/2.3$ HST nominal primary. Apparently, the most extensive use of fast compensators can be expected in the domain of wide-field survey telescopes \\citep{b11}.  \\begin{figure} \\includegraphics[width=84mm]{fig01.eps} \\caption{Verification the HST compensator in autocollimating mode.} \\label{f01} \\end{figure} ",
        "conclusions": "As we see, the proposed three-lens compensator has two core features: \\begin{itemize} \\item There is a position of a point light source, at which light twice passes through all lenses and forms a diffraction image at the source position. It allows us to find small deviations from nominal characteristics not only of the lenses, but also their mutual positions. \\item The radius of curvature of the first compensator's surface is equal in the absolute value to the distance of the compensator from the light source in a control mode. This feature allows us to set correctly both the light source and compensator. \\end{itemize}  It is worthy to mention, as an additional feature, that all the lenses are rather small, so one can choose a homogeneous piece of glass of which lenses will be made. Besides the specified, there are some other features of individual samples. In particular, due to moderate speed of G.A.~Shain primary mirror, $f/3.85$, it was possible to reduce the compensator diameter to $70$~mm and make some of its curvature radii identical.  The autocollimating compensator can be considered as the development of a three-lens design by \\citet{b8} intended for the control of the 6-m BTA primary mirror. In the Puryaev's compensator, the `setting' reflection of light occurs not from the first surface, as in our design, but from the front surface of the second lens; this feature brings an evident uncertainty, because the error in distance of the light source can be balanced by the deviation of geometrical parameters of first two lenses or by the refraction index of the first lens. Eventually, the basic difference of these two schemes is that our design allows the self-test both in the assembled and control states."
    },
    "1402/1402.4298_arXiv.txt": {
        "abstract": " ",
        "introduction": "Chemically peculiar stars of the upper main sequence cover the whole range of spectral types from early B to early F and and do not seem to be constrained to a particular evolutionary stage (Kochukhov \\& Bagnulo, 2006). In addition to their strong, globally strucured magnetic fields, many Ap stars also show prominent spectral line profile variations synchronized to stellar rotation, which, within the framework of the oblique rotator model (introduced by Stibbs (1950), is attributed to oblique magnetic and rotation axes and to the presence of a non-uniform distribution of chemical elements on their surface. With few exceptions such inhomogeneities exist only in the atmospheres of A stars with magnetic fields, demonstrating that these fields play a crucial role in their formation and evolution. The mentioned chemical peculiarities are attributed to the selective diffusion of ions under the competitive action of radiative acceleration and gravitational settling within an atmosphere that is stabilized by the magnetic field (Michaud, 1970). The observed field structures and the chemical spots typically remain stable over many decades.  Modeling of magnetic field effects is extremely challenging, both from the observational and the theoretical point of view, and only recent groundbreaking advances in observational instrumentation, as well as a deeper theoretical understanding of magnetohydrodynamic processes and diffusion in stellar atmospheres (e.g. Alecian et al. 2011) enable us to model stellar fields and their relation to chemical inhomogeneities in more and more detail.  At the same time, the development and application of magnetic Doppler imaging (Kochukhov et al. 2004, Donati et al. 2006, L\\\"uftinger et al. 2010a,b), a method that has become one of the most powerful remote sensing methods in present astrophysics, allows a detailed confrontation of these new models with observations. ",
        "conclusions": "For several decades, DI was applied only to a small number of chemical species, and magnetic fields were analysed with coarse techniques, assuming low-order multipolar magnetic topologies inferred from longitudinal magnetic field measurement. Recent groundbreaking advances in theoretical modeling, observational instrumentation and analysis techniques enabled an enormous development of the (M)DI technique, and we now have reached a stage of maturity, where chemical geometry and magnetic field structures on the surfaces of Ap stars can be determined simultaneously and consistently, and, in the case of full Stokes parameter inversions, without having to rely on any a priory assumptions. These advances have lead to a significant increase of the number of stars mapped. The maps reveal that Ap star surfaces show a pronounced diversity of spot structures, still to be explained. The inclusion of linear Stokes parameters in MDI indicates complex patches of field intensitiy distributions beneath a dominantly dipolar overall field configuration. The small scale structures due to the horizontal magnetic field component possibly influence the aforementioned diversity of chemical patches and an apparent lack of correlation to the overall topology for most of the chemical species. Bayesian Photometric imaging since recently offers an additional asset and might help in increasing the statistics of stars analysed via DI, MDI, and BPI, trying to find an overall picture of the interplay and the relation of diffusion, chemical surface patterns and magnetic field geometries in Ap stars' atmopheres.  \\bigskip {\\it Acknowledgements.} This work was supported by the  Austrian Science Fund (FWF-P19962). The author thanks O. Kochukhov and T. Ryabchikova for useful discussions. Scientific analysis made use orm f the SIMBAD, and VALD databases."
    },
    "1402/1402.4251_arXiv.txt": {
        "abstract": "The entire evolution of the Milky Way, including its mass-assembly and star-formation history, is imprinted onto the chemo-dynamical distribution function of its member stars, $f(x, v, {\\rm [X/H]})$, in the multi-dimensional phase space spanned by position, velocity, and elemental abundance ratios. In particular, the chemo-dynamical distribution functions for low-mass stars (e.g., G- or K-type dwarfs) are precious tracers of the earliest stages of the Milky Way's formation, since their main-sequence lifetimes approach or exceed the age of the universe. A basic tenet of essentially all previous analyses is that the stellar metallicity, usually parametrized as [Fe/H], is conserved over time for main-sequence stars (at least those that have not been polluted due to mass transfer from binary companions). If this holds true, any correlations between metallicity and kinematics for long-lived main-sequence stars of different masses, effective temperatures, or spectral types must strictly be the same, since they reflect the same mass-assembly and star-formation histories. By analyzing a sample of nearby metal-poor halo and thick-disk stars on the main sequence, taken from Data Release 8 of the Sloan Digital Sky Survey, we find that the median metallicity of G-type dwarfs is systematically higher (by about 0.2 dex) than that of K-type dwarfs having the same median rotational velocity about the Galactic center. If it can be confirmed, this finding may invalidate the long-accepted assumption that the atmospheric metallicities of long-lived stars are conserved over time. ",
        "introduction": "Numerical calculations of the nuclear processes that take place in stellar interiors have established that the evolutionary track of a given star is determined primarily by two parameters -- its mass and chemical composition (other parameters, such as stellar rotation and magnetic fields may play roles as well, but they are expected to be smaller). While main-sequence stars are believed to retain the initial values of their mass and composition from birth through the completion of core H burning, it has been conjectured previously that this may not necessarily be the case, depending on their surrounding environment. For example, \\cite{Bondi1952} demonstrated, from theoretical considerations, that a star can accrete gas onto its surface at a rate proportional to $v_{\\rm rel}^{-3}$, where $v_{\\rm rel}$ is the relative velocity between the star and the gas with which it collides. In particular, if initially metal-poor stars accrete metal-rich gas, this would lead to an enhancement of the metallicity in their atmospheres, and thus confound a straightforward interpretation of the chemical evolution of the Galaxy.  \\cite{Yoshii1981} was the first to speculate that Bondi mass accretion onto halo stars may drastically alter their surface metal abundances in the course of formation of the Galactic halo, proposing that more metal-poor halo stars would suffer greater surface metal enhancements than metal-rich stars, due to their shallower surface convective envelopes. He argued that, early in the formation history of the Milky Way, before the emergence of the disk, gaseous material is primarily distributed throughout the halo. During this era, high-density gas clouds move more or less randomly throughout the halo. If halo stars formed from these clouds, it is possible for them to collide at small relative velocity with chemically processed dense gas from previous star-formation episodes in these same clouds, so that mass accretion as well as surface metal enhancement is likely to occur. After the halo gas has contracted to form the disk stellar populations, mass accretion onto halo stars becomes negligible, for two reasons. First, halo stars collide with dense gas only occasionally, when they pass through the disk. Secondly, such collisions occur near the pericenter of the stellar orbits, hence the relative velocity between halo stars and the disk gas is too large to enable efficient mass accretion.  According to the currently favored hierarchical galaxy formation paradigm in a $\\Lambda$-CDM universe, the Galaxy is thought to form through mergers of sub-galactic systems for which the internal velocity dispersion is much smaller than that of the ensemble of these systems that make up the entire halo (see, e.g., \\citealt{Tissera2013}, and references therein). Because the relative velocity between the stars and the gas within these systems is on the order of their internal velocity dispersions, this scenario implicitly involves conditions that are favorable for mass accretion, as suggested years ago by \\cite{Yoshii1981}.  Several previous authors have attempted to evaluate the rate of mass accretion in a cosmological context, and reached similar conclusions, that the surface metal abundances of low-mass, initially metal-poor (or even near zero-metallicity stars) could indeed be enhanced significantly (e.g., \\citealt{Shigeyama2003, Komiya2010}). Other authors have argued against the notion that significant mass accretion onto halo stars could have occurred. For example, \\cite{Frebel2009} calculated the orbits of individual halo stars in order to evaluate the amount of mass that might have been accreted as they penetrate the disk (several times over the age of the Galaxy), and concluded that their surface metal abundances are hardly enhanced.  However, in this work, the amount of mass accreted by individual stars within their natal sub-galactic systems, which may be the dominant source, is not taken into account.  In order to break this stalemate, we propose an observational test to decide whether or not the accretion hypothesis may indeed be supported. Suppose that metal-poor halo G- and K-type dwarfs, with main-sequence lifetimes close to or exceeding the age of the universe, experience no mass accretion and associated surface metallicity enhancement, and that their kinematical properties are independent of spectral type, because they share the same star-formation and mass-assembly histories. Then, any observed correlation between metallicity and kinematics for G-type dwarfs must be identical to that observed for K-type dwarfs. However, if mass accretion and metallicity enhancement did occur, the situation is expected to be quite different. According to theoretical studies of the internal structures of stars, the mass of the surface convective envelope for dwarfs drastically decreases with increasing effective temperature, $T_{\\rm eff}$. This implies that the affect of surface metal enhancement on G-type dwarfs is expected to be much larger than for K-type dwarfs (due to the lack of dilution in their surface layers), even if they had the same initial metal abundance and the same amount of mass accreted. Thus, we can accept or reject the accretion hypothesis by examining whether a spectral-type dependent shift in metallicity for stars with otherwise identical kinematics exists or not.  Note that this test is possible only statistically, based on a large, kinematically unbiased sample of metal-poor main-sequence dwarfs. In this paper, we employ spectroscopic observations of G- and K-type dwarfs in a relatively local region of the halo, taken from Data Release 8 of the Sloan Digital Sky Survey (SDSS DR8, \\citealt{Aihara2011}), to derive the relation between the median rotational velocity, $\\langle V_\\phi \\rangle_{\\rm med}$, and the surface metal abundance, [Fe/H], for stars of different spectral types. If the above mentioned shift in metallicity with respect to kinematics is confirmed, it would be the first observational evidence of the importance of mass accretion onto halo stars. Given the implications of this, and the clear influence on the interpretation of future observations, it is crucial to seek confirmation (or refutation) based on additional studies.  This paper is organized as follows. We describe our sample selection in Section 2, and present our analysis and results in Section 3. The plausibility of our results is discussed in Section 4. Our interpretation of these results is presented in Section 5. Finally, in Section 6 we summarize our results, and suggest future tests of this hypothesis. ",
        "conclusions": "In this paper we have derived the relation of $\\langle V_\\phi \\rangle_{\\rm med}$ (median value of $V_\\phi$) as a function of [Fe/H] for halo G- and K-type dwarfs from SDSS DR8. We find that the run of $\\langle V_\\phi \\rangle_{\\rm med}$ vs. [Fe/H] can be characterized by a boundary metallicity, [Fe/H$]_{\\rm knee}$, below which the nearly non-rotating halo stars dominate, and that [Fe/H$]_{\\rm knee}$ for G-type dwarfs occurs at a higher abundance than that for K-type dwarfs by an offset of $\\delta \\simeq 0.20\\;{\\rm dex}$. This offset is also seen for those sample stars that have large vertical motion. Below we consider the implications of this result.   \\subsection{Evidence of Metal Accretion onto Main-Sequence Halo Stars} \\label{evidence}  We interpret this non-zero value of $\\delta \\simeq 0.20\\;{\\rm dex}$ between the characteristic metallicities of halo G- and K-type dwarfs to be the first tentative evidence that halo stars have been {\\it externally} polluted by the accretion of metal-enriched gas from their natal clouds. The effect is more noticeable for G-type dwarfs, whose convective envelopes are shallower than those of K-type dwarfs, since they will not have diluted accreted material as fully.  If this interpretation is correct, we expect an even larger offset in metallicity for red giants (e.g., K giants), which possess deeper surface convective envelopes than K-type dwarfs \\citep{Yoshii1981}. That is, the knee of the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation for K giants, if plotted as in Figure \\ref{fig 1 pm}, should be located to the left of that for K dwarfs. Although there is no large publicly available database of nearby K giants, other types of post-main-sequence stars might serve as alternatives, since the surface heavy element abundance is expected to remain unchanged when red giants evolve into either horizontal-branch stars or RR Lyrae stars. During the course of this evolution, red giants lose mass from their surface convective envelopes. The total amount of mass loss is expected to be $\\simeq0.2\\; M_\\odot$ for a red giant of initially $0.8\\;M_\\odot$ (see \\citealt{Yoshii1981}, and references therein), which is smaller than the total mass of the surface convective envelope ($\\simeq0.3\\;M_\\odot$) in the red-giant stage \\citep{Sweigart1978}.  Thus, instead of red giants, we have explored the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation for 290 blue horizontal-branch (BHB) stars with $T_{\\rm eff} > 7500\\;{\\rm K}$, $\\log g <3.8$, located at $2\\; {\\rm kpc}<|z|<5\\;{\\rm kpc}$, $5\\;{\\rm kpc}<R<20\\;{\\rm kpc}$, taken from \\cite{Xue2011}. Here, $R$ and $|z|$ are the Galactocentric distance projected onto the Galactic disk plane and the distance above or below the disk plane, respectively. The cuts in $T_{\\rm eff}$ and $\\log \\; g$ are adopted following the selection criteria proposed by R. Santucci et al. (in preparation).  In Figure \\ref{fig 3 BHB}, we show the distribution of BHB stars in the $V_\\phi$-[Fe/H] space, along with the median value $\\langle V_\\phi \\rangle_{\\rm med}$ for the binned sample. From inspection of this figure, the $\\langle V_\\phi \\rangle_{\\rm med}$ of BHB stars is nearly constant (consistent with 0 ${\\rm km\\;s^{-1}}$) at [Fe/H] $<$ [Fe/H$]_{\\rm knee} \\simeq -1.7$.\\footnote{ The plateau value of $\\langle V_\\phi \\rangle_{\\rm med}$ at [Fe/H] $<$ [Fe/H$]_{\\rm knee}$ for the BHB sample ($\\simeq 0\\;{\\rm km\\;s^{-1}}$) is lower than that for G/K-type dwarfs ($\\simeq 50\\;{\\rm km\\;s^{-1}}$; see Figure \\ref{fig 1 pm}). This discrepancy seems to arise from the different spatial regions covered by these samples. In fact, \\cite{CB2000} find that the mean rotational velocity $\\langle V_\\phi \\rangle$ for nearby metal-poor stars with [Fe/H] $< -1.5$ is higher for those sample stars with smaller $|z|$ [see their Figure 3(a)], which is consistent with our discrepancy. } This suggests that inner-halo BHB stars dominate over metal-weak thick-disk BHB stars below [Fe/H] $=$ [Fe/H$]_{\\rm knee}$ at $2\\; {\\rm kpc}<|z|<5\\; {\\rm kpc}$. Noting that a lower value of [Fe/H$]_{\\rm knee}$ is expected for a sample of stars with lower $|z|$ (due to the larger fraction of thick-disk stars in the sample), we expect that [Fe/H$]_{\\rm knee}$ for BHB stars would be even lower than $-1.7$ if we could obtain a BHB sample with distances in the range 0.84 kpc $< d <$ 1.64 kpc, as was used above for the G- and K-dwarf samples. Since the positions of [Fe/H$]_{\\rm knee}$ for the G and K-type dwarfs are $-1.4$ and $-1.6$, respectively, in this distance range (see Figure \\ref{fig 1 pm}), it follows that [Fe/H$]_{\\rm knee}$ for BHB stars should be lower, by {\\it at least} 0.3 dex and 0.1 dex, respectively, if these samples could be fairly compared. This may indicate that the BHB stars are even less affected by surface metal pollution than K-type dwarfs. The sample of 290 BHB stars is not sufficiently large to be certain of this effect, but the present result does serve to support the accretion hypothesis for halo stars. Surveys such as LAMOST and Gaia will provide larger kinematically unbiased samples of low-mass dwarfs and red giants. At that stage, we will be able to compare the chemo-dynamical correlations of these stars more rigorously, and test the metal accretion hypothesis thoroughly.   \\subsection{Where did the Metal Accretion Take Place?}  As mentioned in the Introduction, the efficiency of metal accretion in halo stars is proportional to $v_{\\rm rel}^{-3}$, where $v_{\\rm rel}$ is the relative velocity of the star and the colliding gas. Therefore, metal accretion onto halo stars may or may not be important, depending on the nature of the environments in which it could take place. Recent observations \\citep{Carollo2007, Carollo2010, Beers2012} and numerical simulations \\citep{Font2011, McCarthy2012, Tissera2013} suggest that the stellar halo of the Milky Way comprises two distinct components with different origins, which are often referred to as the inner- and outer-halo populations. In this picture, inner-halo stars formed from relatively more massive sub-galactic systems ($10^{9-10} M_\\odot$) in the main progenitors of the Milky Way, while much smaller sub-galactic systems ($10^{7-8} M_\\odot$ or less), similar to lower-mass dwarf-like galaxies, were disrupted to contribute the bulk of the low-metallicity outer-halo stars. In the progenitor systems of the inner-halo stars, the internal velocity dispersion is expected to be $\\sim 30\\;{\\rm km\\; s^{-1}}$ \\citep{Tissera2013}, so that the above-mentioned $v_{\\rm rel}$ may have been too large for inner-halo stars to experience efficient metal accretion. On the other hand, the internal velocity dispersion of ultra-faint dwarf galaxies -- whose constituent stars are typically as metal-poor as [Fe/H] $\\lesssim -2.0$, similar to nearby outer-halo stars -- is $\\sim 5\\;{\\rm km\\; s^{-1}}$ (\\citealt{Simon2007, Walker2009}; see also \\citealt{YoshiiArimoto1987}), so that $v_{\\rm rel}$ may have been small enough to enable efficient metal accretion onto constituent stars. Therefore, in this inner/outer halo picture, the outer-halo stars are much more likely to have experienced metal accretion. In other words, if the dual halo picture is correct, we expect that inner-halo dominated sample of G- and K-type dwarfs would exhibit little or no offset in the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation, while the outer-halo dominated sample of G- and K-type dwarfs would show noticeable offsets.  Observationally, it is suggested that the fraction of outer-halo stars increases as $z_{\\rm max}$ increases \\citep{Carollo2010}. Thus, if we divide each of our G- and K-dwarf samples into two subsamples with respect to $z_{\\rm max}$, the fraction of outer-halo stars is expected to be larger in the higher-$z_{\\rm max}$ subsample than in the lower-$z_{\\rm max}$ one. It follows that we expect a clearer offset in the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation for the higher-$z_{\\rm max}$ subsample. This expectation is indeed consistent with the results in Figure \\ref{fig 2 zmax}, in which we observe a clearer offset in the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation for our sample stars with $z_{\\rm max} > 3 \\; {\\rm kpc}$, and not for those stars with $z_{\\rm max} < 3 \\; {\\rm kpc}$.\\footnote{ Our boundary at $z_{\\rm max} =$ 3 kpc is smaller than that adopted in \\cite{Carollo2010}, so that we have a sufficient number of stars in the higher-$z_{\\rm max}$ subsample. However, this difference is not crucially important, since our aim is to examine two subsamples with different fractions of inner/outer-halo stars. Our higher-$z_{\\rm max}$ subsample includes numerous stars with $z_{\\rm max}$ as large as 10-50 kpc. } These results can be understood if a dual inner/outer halo applies to the Milky Way, as originally suggested by \\cite{Carollo2007}, and metal accretion is efficient only for stars associated with the progenitors of the outer halo.  We note here that typical halo stars with $z_{\\rm max}>$ 3 kpc would penetrate through the Galactic disk with a velocity of $v_{\\rm rel} > 100\\;{\\rm km\\;s^{-1}}$ relative to the gas clouds in the disk. Therefore, metal accretion when the halo stars penetrate the Galactic disk plane accounts for only a negligible fraction of the total amount of metals accreted onto halo stars, supporting the view of \\cite{Frebel2009}. However, the fact that we see a clear offset in panel (b) of Figure 3 suggests that the outer-halo stars are likely to have experienced metal accretion within their progenitor systems, which were {\\it later} disrupted.   \\subsection{Impact on the Metallicity Distribution Function in the Stellar Halo System}  One obvious impact of our results is that the metallicity distribution function (MDF) of halo main-sequence stars may need to be re-examined. In the hierarchical galaxy formation scenario, halo stars form in sub-galactic systems that are later disrupted by tidal interaction with the Galaxy, hence halo stars originating from these disrupted systems are likely to have experienced metal accretion within these systems \\citep{Shigeyama2003, Suda2004, Komiya2010}. As a consequence, the shape of the MDF of halo stars is expected to be skewed toward higher [Fe/H] when compared to their original MDF, because the surface metallicity is enhanced for stars with lower [Fe/H] due to their shallow surface convective envelopes, while it is less enhanced for stars with higher [Fe/H] due to their deeper surface convective envelopes \\citep{Mengel1979}.  In this respect, it is intriguing to note the apparent discrepancy between the MDFs of main-sequence turn-off (MSTO) stars and BHB stars in the outer halo region of the Milky Way. \\cite{Sesar2011} show that the median metallicity of MSTO stars at Galactocentric distances in the range of $10\\;{\\rm kpc}<r<30\\;{\\rm kpc}$ is [Fe/H]$\\simeq -1.5$, almost independent of $r$.\\footnote{ See panels W3 and W4 in Figure 12 of \\cite{Sesar2011}, in which the influence from known substructures in the halo is not significant.} On the other hand, \\cite{Beers2012} show that the median metallicity of BHB stars at $10\\;{\\rm kpc}<r<40\\;{\\rm kpc}$ is [Fe/H]$\\simeq-2.0$, almost independent of $r$.\\footnote{ See the right-hand panels in Figure 15 of \\cite{Beers2012}, in which the contamination from the Sagittarius dwarf galaxy is minimized, and Figure 4 of \\cite{Carollo2007}. } This discrepancy can be well explained by the accretion hypothesis. MSTO and BHB stars are of similar age, yielding no systematic difference in the mass accreted in their main-sequence stage. Thus, the surface metallicity of BHB stars is much less enhanced than that of MSTO stars, because the red giants that were progenitors of the stars presently on the BHB had very deep surface convective envelopes compared to MSTO stars (see section \\ref{evidence}).  Furthermore, noting that the surface convective envelopes of MSTO stars are shallower than those of G-type dwarfs, it is also interesting to note that the apparent offset in the median metallicity of the MDF for MSTO and BHB stars ([Fe/H]$_{\\rm median, MSTO} - $ [Fe/H]$_{\\rm median, BHB} \\simeq 0.5 \\;{\\rm dex}$) is consistent with the offset in the break metallicity ([Fe/H$]_{\\rm knee, G} - $[Fe/H$]_{\\rm knee, BHB} \\gtrsim 0.3$ dex) of the $\\langle V_\\phi \\rangle_{\\rm med}$-[Fe/H] relation for G-type dwarfs and BHB stars (see section \\ref{evidence}). Our interpretation of this consistency in terms of the accretion hypothesis may be tested based on a much larger sample of BHB stars or red giants that should be available in the near future.  Another intriguing aspect of the observed halo MDF is the possible existence of a cutoff metallicity at [Fe/H]$_{\\rm cutoff}\\; \\simeq -4$, below which the MDF shows a sharp decline (\\citealt{Schorck2009, Li2010}; see also \\citealt{Yong2013}). There are not yet a sufficient number of stars known with [Fe/H] $< -3.5$, let alone with [Fe/H] $< -4.0$, to evaluate whether the cutoff is real, or simply the result of small number statistics in the ultra metal-poor regime. Future observations of G/K-type dwarfs might well detect a spectral-type dependence of this cutoff metallicity, in the sense that [Fe/H]$_{\\rm cutoff, K} < $ [Fe/H]$_{\\rm cutoff, G}$, if metal accretion and subsequent surface metal enhancement does indeed take place among metal-poor halo stars.   \\subsection{On the Case of $\\omega$ Centauri}    Some globular clusters appear to be ideal places to test the metal accretion hypothesis, if they are regarded as closed systems and the constituent stars are coeval, or nearly so. \\cite{Stanford2007} investigate the metallicities of stars in the globular cluster $\\omega$ Cen, and find that the peak metallicity of near-turn-off stars is systematically higher than that of red giant branch (RGB) stars, by 0.01-0.05 dex (see their Figure 2). This offset is smaller than the uncertainty in the derived [Fe/H] for their sample (0.15-0.20 dex). However, taking into account that the surface convective envelope is deeper for RGB stars than near-turn-off stars, the sense of the reported gap is in agreement with the prediction from the metal accretion hypothesis.   If the reported gap (0.01-0.05 dex) between near-turn-off stars and RGB stars is due to metal accretion, we expect that the typical metallicities of G- and K-type dwarfs in $\\omega$ Cen are different from each other by less than 0.05 dex, because of the milder difference in the depth of surface convective envelope between G- and K-type dwarfs. If we take into account the observed offset of $\\simeq 0.20$ dex in Figure \\ref{fig 2 zmax}(b), the metal accretion in typical progenitor systems of the Milky Way outer halo is more significant than that in $\\omega$ Cen. Noting that the efficiency of metal accretion is proportional to $v_{\\rm rel}^{-3}$, the seemingly low efficiency within $\\omega$ Cen indicates that typical progenitor systems of the outer halo must have smaller internal velocity dispersions than that of $\\omega$ Cen. It is suggested that such progenitor systems are similar to the currently observed ultra-faint dwarf galaxies (with central velocity dispersion $\\sim 5 \\; {\\rm km\\;s^{-1}}$), and that the contribution from large and compact globular clusters (such as $\\omega$ Cen, with central velocity dispersion $\\simeq 17 \\; {\\rm km\\;s^{-1}}$; \\citealt{Sollima2009}) to the outer halo is not significant.     \\subsection{Implication for the First Stars}  If primordial (zero-metal) stars from the very first generation (with main-sequence masses below $0.8 M_\\odot$) were able to form, they could survive until today, and possibly be observed.  Here we consider this possibility from the perspective of the accretion hypothesis.  The mass range of the first stars has been long debated, but some nucleosynthesis constraints have recently been placed by the discovery of four hyper metal-poor (HMP; technically, Fe-poor) stars with [Fe/H]$\\lesssim-5$ (HE~01072-5240: \\citealt{Christlieb2002}; HE~1327-2326: \\citealt{Frebel2005}; HE~0557-4840: \\citealt{Norris2007}; SDSS~J102915+172927: \\citealt{Caffau2012}). Three of the four stars (excluding SDSS~J102915+172927) exhibit elemental abundance patterns that are significantly enhanced with carbon (C) and nitrogen (N), and can be well-explained if they are second-generation stars, formed out of an interstellar medium that was chemically enriched by supernova ejecta from $25 M_\\odot$ primordial stars \\citep{Umeda2003, Iwamoto2005, Nomoto2013}.  However, if we allow for the metal accretion hypothesis, as suggested to apply in this paper, it is possible to regard the above-mentioned HMP stars as surviving primordial stars whose surfaces were polluted by the accretion of the supernova ejecta of $25 M_\\odot$ primordial stars \\citep{Shigeyama2003}. Since the mass of these stars is as small as 0.8$M_\\odot$, our interpretation has an implicit requirement for the mass of primordial stars to range from less than $1 M_\\odot$ up to more than a few tens of $M_\\odot$, as predicted by some authors \\citep{Yoshii1986, Nakamura2001}. In this respect, it is worth mentioning that SDSS~J102915+172927 is an 0.7$M_\\odot$ star with no significant C and N enhancement. The mass fraction of heavy elements derived for this star, $Z<7.4\\times10^{-7}$, is the lowest known to date, and heavy-element cooling in metal-poor gas with $Z<Z_{\\rm crit} = 10^{-5} - 10^{-6}$ is superseded by purely atomic or molecular hydrogen or zero-metal cooling \\citep{Silk1977, Yoshii1980, Bromm2001, Smith2009, Safranek2010}. Therefore, the discovery of a 0.7$M_\\odot$ star with $Z < Z_{\\rm crit}$ indicates that this star could indeed have formed by zero-metal cooling. In other words, the zero-metal cooling process turns out to be able to produce primordial stars with the masses below $1 M_\\odot$.  We propose that at least some of the HMP stars could be surviving primordial stars that have experienced surface metal pollution after they formed. Given that the environments in which such pollution can occur is limited, we expect that even more HMP metal-poor stars similar to SDSS~J102915+172927, with $Z \\ll Z_{\\rm crit}$, may well be discovered."
    },
    "1402/1402.6254_arXiv.txt": {
        "abstract": "\\noindent {\\ftn \\bf\\scshape Abstract:} We explore the upper bound on the tensor-to-scalar ratio $r$ in supersymmetric (F-term) hybrid inflation models with the gauge symmetry breaking scale set equal to the value $2.86\\cdot10^{16}~{\\rm GeV}$, as dictated by the unification of the MSSM gauge couplings. We employ a unique renormalizable superpotential and a quasi-canonical K\\\"ahler potential, and the scalar spectral index $n_s$ is required to lie within the two-sigma interval from the central value found by the Planck satellite. In a sizable region of the parameter space the potential along the inflationary trajectory is a monotonically increasing function of the inflaton, and for this case, $r\\lesssim2.9\\cdot10^{-4}$, while the spectral index running, $|d\\ns/d\\ln k|$, can be as large as $0.01$. Ignoring higher order terms which ensure the boundedness of the potential for large values of the inflaton, the upper bound on $r$ is significantly larger, of order $0.01$, for subplanckian values of the inflaton, and $|d\\ns/d\\ln k|\\simeq0.006$. \\\\ \\\\ {\\scriptsize {\\sf PACs numbers: 98.80.Cq, 12.60.Jv}   } ",
        "introduction": "\\emph{Supersymmetric} ({\\small\\sf SUSY}) hybrid inflation based on F-terms, also referred to as \\emph{F-term hybrid inflation} ({\\sf \\small FHI}), is one of the simplest and well-motivated inflationary models \\cite{susyhybrid,hybrid}. It is tied to a renormalizable superpotential uniquely determined by a global $U(1)$ R-symmetry, does not require fine tuned parameters and it can be naturally followed by the breaking of a \\emph{Grand Unified Theory} ({\\sf \\small GUT}) gauge symmetry, such as $G_{B-L}= G_{\\rm SM}\\times U(1)_{B-L}$ \\cite{bl}, where ${G_{\\rm SM}}= SU(3)_{\\rm C}\\times SU(2)_{\\rm L}\\times U(1)_{Y}$ is the gauge group of the \\emph{Standard Model} ({\\small\\sf SM}), $\\Glr=SU(3)_{\\rm C}\\times SU(2)_{\\rm L} \\times SU(2)_{\\rm R} \\times U(1)_{B-L}$ \\cite{dvali, vlachos}, and flipped $SU(5)$ \\cite{barr, ant, davis, flipped}, with gauge symmetry $\\Gfl=SU(5)\\times U(1)_X$. Let us clarify, in passing, that the term ``GUT'' is used in the sense of the gauge coupling unification within \\emph{Minimal SUSY SM} ({\\sf \\small MSSM}), although the aforementioned gauge groups are not simple. Such models can arise from string compactifications, see for e.g. \\cref{tracas,ant}. The embedding of the simplest model of FHI within a higher gauge group may suffer from the production of cosmic defects which can be evaded, though, in the cases of smooth \\cite{smooth} or shifted \\cite{shifted} FHI.   In the simplest implementation of FHI \\cite{susyhybrid}, we should note that the potential along the inflationary track is completely flat at tree level. The inclusion of \\emph{radiative corrections} ({\\sf \\small  RCs}) \\cite{susyhybrid} produce a slope which is needed to drive inflaton towards the SUSY vacuum. In this approximation the predicted scalar spectral index $n_s\\simeq0.98$, is in slight conflict with the latest WMAP \\cite{wmap} and PLANCK \\cite{plin} data based on the standard power-law cosmological model with \\emph{Cold Dark Matter and a cosmological constant} ({\\sf \\small $\\Lambda$CDM}). Furthermore, the gauge symmetry breaking scale $M$ turns out to be close to (but certainly lower than) its SUSY value, $\\Mgut\\simeq2.86\\cdot10^{16}~{\\rm GeV}$.  A more complete treatment which incorporates \\emph{supergravity} ({\\sf \\small SUGRA}) corrections \\cite{senoguz} with canonical (minimal) \\Kap, as well as an important soft SUSY breaking term \\cite{sstad,mfhi}, has been shown to yield values for $\\ns$ that are fully compatible with the data \\cite{wmap,plin}, with $M$ in this case somewhat lower than the one obtained in \\cref{susyhybrid}. A reduction of $M$ is certainly welcome if FHI is followed by the breaking of an abelian gauge symmetry, since it helps to reconcile $M$ with the bound \\cite{plcs} placed on it by the non-observation of cosmic strings \\cite{jp, gmb, mfhi, buch}.  The minimal FHI scenario described above, while perfectly consistent with the current observations, requires some modification if one desires to incorporate values of $M$ that are comparable or equal to $\\Mgut$. This is indispensable in cases where $\\Ggut$ includes non abelian factors besides $\\Gsm$, which are expected to disturb the successful gauge coupling unification within MSSM. In this letter, we would like to emphasize that the observationally favored values (close to 0.96) for $\\ns$ with $M$ equal to the SUSY GUT scale can be readily achieved within FHI by invoking a specific type of non-minimal \\Kap, first proposed in \\cref{rlarge}. In particular, a convenient choice of the next-to-minimal and the next-to-next-to-minimal term of the adopted \\Ka generates \\cite{rlarge,hinova,alp} a positive mass (quadratic) term for the inflaton and a sizeable negative quartic term which assist us to establish FHI of hilltop type \\cite{lofti} in most of the allowed parameter space of the model. Our objectives can also be achieved in smaller regions of the allowed parameter space even with monotonic inflationary potential and therefore complications related to the initial conditions of FHI can be safely eluded. Acceptable $\\ns$ values within this set-up are accompanied with an enhancement of the running of $\\ns$, $\\as$, and the scalar-to-tensor ratio, $r$, which reach, thereby, their maximal possible values within FHI if we take into account that $M$'s larger than $\\Mgut$ are certainly less plausible. Note, in passing, that the reduction of $\\ns$ by generating a negative mass (quadratic) term for the inflaton, as done in \\cref{gpp}, is not suitable for our purposes since $M$ remains well below $\\Mgut$.   Below, we briefly review in \\Sref{fhi} the basics of FHI when it is embedded in nonminimal SUGRA and recall in \\Sref{fhi3} the observational and theoretical constraints imposed on our model. In \\Sref{res} we exhibit our updated results, and our conclusions are summarized in \\Sref{con}. ",
        "conclusions": "\\label{con}   Inspired by the recently released results by the PLANCK collaboration on the inflationary observables, we have reviewed and updated the nonminimal version of SUSY hybrid inflation arising from F-terms, also referred to as FHI. In our formulation, FHI is based on a unique renormalizable superpotential, employs an quasi-canonical \\Ka\\ and is followed by the spontaneous breaking at $\\Mgut$ of a GUT symmetry which is taken to be $\\Glr$ or $\\Gfl$. As suggested first in \\cref{rlarge} and further exemplified in \\cref{hinova,alp}, $\\ns$ values close to $0.96$ in conjunction with the fulfilment of \\Eref{Mgut} can be accommodated by considering an expansion of the \\Ka\\ -- see \\Eref{K} -- up to twelfth order in powers of the various fields with suitable choice of signs for the coefficients $\\ks$ and $\\kss$.  Fixing $n_{\\rm s}$ at its central value, we obtain $\\{7.8\\cdot10^{-2}\\}~1.57\\lesssim{\\kp}/{10^{-2}}\\lesssim2.2$ with $\\{2\\}~4.2\\lesssim{-\\ks}/{10^{-2}}\\lesssim7.2$ and $0.72~\\lesssim {\\kss}\\lesssim0.79~\\{10\\}$, while $|\\as|$ and $r$ assume the values $(\\{0.1\\}~0.45-1)\\cdot10^{-2}$ and $\\lf\\{3.5\\cdot10^{-3}\\}~1.4-1.9\\rg\\cdot10^{-4}$ respectively -- recall that the limiting values in the curly brackets are achieved without imposing the monotonicity of $\\Vhi$. With a non-monotonic $\\Vhi$, $\\Dex$ ranges between $16$ and $30\\%$.  It is gratifying that there is a sizable portion of the allowed parameter space where $\\Vhi$ remains a monotonically increasing function of $\\sg$; thus, unnatural restrictions on the initial conditions for inflation due to the appearance of a maximum and a minimum of $\\Vhi$ can be avoided. On the other hand, if we do not insist on the boundedness of $\\Vhi$, $\\kp$ reaches $0.1$ with $\\ks=-0.046$ and $\\kss=0.4$ with the resulting $\\as$ and $r$ being both $0.006$, that is close to $0.01$. Finally FHI can be followed by a successful scenario of non-thermal leptogenesis \\cite{lept} for both $\\Ggut$'s considered here -- cf. \\cref{flipped,alp}."
    },
    "1402/1402.6306_arXiv.txt": {
        "abstract": "The theoretical analysis of many problems in physics, astronomy and applied mathematics requires an efficient numerical exploration of multimodal parameter spaces that exhibit broken ergodicity. Monte Carlo methods are widely used to deal with these classes of problems, but such simulations suffer from a ubiquitous sampling problem: the probability of sampling a particular state is proportional to its entropic weight. Devising an algorithm capable of sampling efficiently the full phase space is a long-standing problem. Here we report a new hybrid method for the exploration of multimodal parameter spaces exhibiting broken ergodicity. Superposition enhanced nested sampling (SENS) combines the strengths of global optimization with the unbiased/athermal sampling of nested sampling, greatly enhancing its efficiency with no additional parameters. We report extensive tests of this new approach for atomic clusters that are known to have energy landscapes for which conventional sampling schemes suffer from broken ergodicity. We also introduce a novel parallelization algorithm for nested sampling. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.4135_arXiv.txt": {
        "abstract": "We explore the internal structure of red sequence galaxies in the Coma cluster across a wide range of luminosities ($-17$\\,$>$\\,$M_g$\\,$>$\\,$-22$) and cluster-centric radii ($0$\\,$<$\\,$r_{\\rm{cluster}}$\\,$<$\\,1.3 $r_{200}$). We present the 2D bulge-disc decomposition of galaxies in deep Canada-France-Hawaii Telescope $u,g,i$ imaging using GALFIT. Rigorous filtering is applied to identify an analysis sample of 200 galaxies which are well described by an `archetypal' S0 structure (central bulge + outer disc). We consider internal bulge and/or disc colour gradients by allowing component sizes to vary between bands. Gradients are required for $30\\%$ of analysis sample galaxies. Bulge half-light radii are found to be uncorrelated with galaxy luminosity ($R_e \\sim 1$ kpc, $n\\sim2$) for all but the brightest galaxies ($M_g<-20.5$). The S0 discs are brighter (at fixed size, or smaller at fixed luminosity) than those of star-forming spirals. A similar colour-magnitude relation is found for both bulges and discs. The global red sequence for S0s in Coma hence results from a combination of both component trends. We measure an average bulge $-$ disc colour difference of $0.09\\pm0.01$ mag in $g-i$, and $0.16\\pm0.01$ mag in $u-g$. Using simple stellar population models, bulges are either $\\sim2$-$3\\times$ older, or $\\sim2\\times$ more metal-rich than discs. The trend towards bluer global S0 colours observed further from Coma's core is driven by a significant correlation in disc colour with cluster-centric radius. An equivalent trend is detected in bulge colours at a marginal significance level. Our results therefore favour environment-mediated mechanisms of disc fading as the dominant factor in S0 formation. ",
        "introduction": "Within galaxy clusters, strong radial trends are observed in galaxy morphology (the `morphology-density' relation, \\citealp{Dressler1980,Dressler1997}) and star formation rate (the `colour-density' relation, \\citealp{Balogh2000a,Lewis2002,Balogh2004,Hogg2004}). This builds a picture in which a galaxy is driven from late (e.g. spiral) to early (e.g. E/S0) morphology as it is accreted onto a cluster due to environment-mediated processes. The current position of a galaxy within a cluster is correlated with its infall time \\citep{Gao2004,Smith2012,DeLucia2012,Taranu2012} and velocity \\citep{Oman2013}. Therefore, variations of galaxy structure and colour with cluster-centric radius offer insight into the physical mechanisms which acted during cluster infall, precipitating morphological evolution.  Structural and kinematical similarities between S0s and spiral galaxies \\citep{vdBergh1976,Aragon2006,Bekki2011} imply a evolutionary scenario in which spiral-like progenitors are transformed into early-type galaxies (ETGs). Such transformations would require the erasure of prominent spiral pattern, truncation of the galaxy's star formation, and enhancement of the galaxy's bulge-to-total ratio (B/T $=f_B/(f_B+f_D)$; \\citealp{Dressler1980}). Devoid of fuel for continued star formation, the galaxy would then fade and redden as its stellar populations age, moving from the `blue cloud' to the well-known `red sequence'. The resulting `red and dead' descendant galaxy would be bulge-dominated, with a smooth, corotating disc. The environment-mediated mechanisms proposed to drive this morphological transformation can be broadly categorised as either `bulge-enhancing' or `disc-fading'.  In a `bulge enhancement' scenario, the galaxy's cold gas is rapidly consumed in strong bursts of centrally-concentrated star formation. Additionally, driving the gas towards the galaxy's centre may feed a supermassive black hole, quenching star formation through active galactic nucleus (AGN) feedback \\citep{Silk1998,Schawinski2006,Schawinski2007,Brodwin2013}. Mergers (in groups or the field) and galaxy harassment (in clusters, \\citealp{Moore1996}) are the most widely studied mechanisms of this type, both of which show the potential to form S0s (\\citealp{Bekki1998,Governato2009,Moore1998,Moore1999}). Stellar discs of galaxies are typically disrupted by these processes, however, requiring specific environmental conditions for disc survival \\citep{Hopkins2009} or a long period of disc regrowth \\citep{Kannappan2009}.  Alternatively, a `disc fading' scenario involves either the direct stripping of cold gas from the galaxy's disc (e.g. due to ram pressure \\citealp{Gunn1972,Quilis2000}), or removal of its hot halo gas reservoir over a long period of strangulation \\citep{Larson1980,Balogh2000a}. These mechanisms act preferentially on gas, causing little or no disruption to the galaxy's stellar disc, but act on different timescales. Additionally, disc fading mechanisms do not act in low-density environments, suggesting that transformation would have had to occur since accretion into a group environment.  Recent work studying galaxy kinematics has indicated that most ETGs are rapidly rotating. In this paradigm, elliptical and lenticular galaxies are part of a continuous sequence of rotating, quiescent galaxies with specific angular momentum increasing in magnitude from E to S0a-c \\citep{Emsellem2011,Cappellari2011}. Hence, the S0 morphology encompasses multiple distinct classes of galaxy, differing in luminosity and evolutionary pathway \\citep{vdBergh1990,vdBergh2009a}. Conversely, the distinction between S0s and the most elongated (E7) elliptical galaxies is purely due to observational bias \\citep{vdBergh2009b}. If so, with sufficient signal-to-noise (S/N), discs should be detectable in most elliptical galaxies. Such multicomponent structures for ETGs are consistent with the predictions of photometric modelling (e.g. \\citealp{Rix1990}).  Measurements of broad-band colours can be used to probe a galaxy's star formation history. Using stellar population synthesis models, optical colours can constrain the ages and metallicities of the observed stellar populations, albeit with degeneracies due to the common bias towards redder colours caused by both properties. Many photometric studies of galaxy evolution only investigated {\\emph{ global}} colours (e.g. \\citealp{Gavazzi2010}). While straightforward, this approach glosses over the rich variation of stellar population properties within and between component structures. Alternatively, bulge-disc decomposition of imaging data (e.g. \\citealp{Hudson2010}; hereafter H10, \\citealp{Simard2011,Lackner2012}) separates the photometric contributions of the galaxy bulge and disc, and thus enables stellar population properties of these structures to be measured separately. An age difference between the bulge and disc would constrain the transformation mechanisms that drove the galaxy's evolution in the past.  The well-studied Coma cluster (Abell 1656) possesses one of the richest ETG populations in the local universe. As such, Coma is an ideal laboratory for studying the properties and evolution of ETGs (e.g. \\citealp{Lucey1991,BowLucEll,Jorgensen1999}). In addition, Coma encompasses a wide range of local environment conditions, allowing in-depth investigation of radial trends of environment-mediated processes \\citep{Gavazzi1989,Guzman1992,Carter2008,Gavazzi2010,Smith2012,Cappellari2013,Rawle2013,Lansbury2014}.  In this work, we investigate the bulges and discs of red sequence galaxies in the Coma cluster within an absolute magnitude range $-22<M_g<-17$. This yields a rich initial catalogue of ETGs for analysis, over a wide range of local environment densities.  Rather than forcing a bulge + disc morphology on all galaxies regardless of model suitability, we focus exclusively on a sample of galaxies with S0-like structural morphologies (i.e. central bulge + outer disc; described as ``Classic\"/Type 1 in \\citet{Allen2006}, hereafter `archetypal'). The cherry-picked nature of this sample introduces selection biases as only symmetric galaxies with idealised morphologies are included. However, our aim here is to examine galaxies expected to be well-fit by a bulge + disc model, rather than address a complete sample.  Using high quality Canada-France Hawaii Telescope (CFHT) optical imaging, we perform bulge-disc decompositions with a flexible ($n$ free) \\sersic + exponential model using GALFIT \\citep{GALFIT}. As this data is significantly deeper than the Sloan Digital Sky Survey (SDSS), with better resolution, our decomposition results achieve a greater level of precision than can be achieved using SDSS imaging. As such, we are able to probe the poorly-investigated faint end of Coma's red sequence with far greater reliability.  We discuss multi-band fitting (i.e. fitting using constraints derived from fits to another photometric band) based on the results of two alternative methods, differing in their interpretation of galaxy colour gradients: i) where gradients are purely due to a difference in colour between bulges and discs. ii) where gradients are a combination of colour separation and internal component gradients. This builds on earlier multi-band studies (e.g. \\citealp{Simard2002,Gadotti2009}), providing the first detailed investigation of internal component colour gradients in the context of 2-component \\sersic fits in 2D (although multiwavelength 1-component \\sersic fitting is discussed at length in \\citealp{Haussler2013}).  We investigate the origins and variation of colour-magnitude trends in ETGs in Coma as a means of probing the mechanisms that influenced their evolution history. We address three main questions: do the bulges and discs of ETGs follow a common red sequence slope?; how separated (in colour) are the stellar populations of these components?; in what way do the observed colour distributions vary during cluster infall?  The outline of the paper is as follows: An overview of the CFHT imaging, and description of the initial galaxy sample is presented in $\\S2$. Details of the bulge-disc decomposition carried out using GALFIT are presented in  $\\S3$, including discussion of sample filtering. In $\\S4$, we present the structural, colour-magnitude (including fitting performed both with and without internal gradients), and colour-radius results. We present a discussion of the reported results in $\\S5$, summarise our findings in $\\S6$. A detailed catalogue of fitting measurements are presented in Tables \\ref{result_1} and \\ref{result_2} of Appendix \\ref{Cat}.  Throughout this paper, we use the WMAP7 cosmology: $H_0=70.4\\rm{km}\\rm{s}^{-1}\\rm{Mpc}^{-1}$ (i.e. $h_{70} = 1.01$),$\\Omega_m=0.272$ and $\\Omega_{\\Lambda}=0.728$ \\citep{WMAP7}. Using $z_{\\rm{CMB}}\\rm{(Coma)}=0.024$, the luminosity distance for the Coma cluster is 104.1 Mpc, and the distance modulus, $m - M = 35.09$. At this distance, $1\\minute$ corresponds to 28.9 kpc. Taking a value for velocity dispersion of $\\sigma_{{\\rm Coma}} = 1008$ km s$^{-1}$ \\citep{Struble1999} and virial mass, $M_{200} = 5.1\\times10^{14} h^{-1}_{70}$\\Msol \\citep{Gavazzi2009}, the virial radius, $r_{200}$, for Coma is 2.2 Mpc ($\\sim75\\minute$). ",
        "conclusions": ""
    },
    "1402/1402.3009_arXiv.txt": {
        "abstract": "A non-singular bouncing cosmology is generically obtained in loop quantum cosmology due to non-perturbative quantum gravity effects. A similar picture can be achieved in standard general relativity in the presence of a scalar field with a non-standard kinetic term such that at high energy densities the field evolves into a ghost condensate and causes a non-singular bounce. During the bouncing phase, the perturbations can be stabilized by introducing a Horndeski operator. Taking the matter content to be a dust field and an ekpyrotic scalar field, we compare the dynamics in loop quantum cosmology and in a non-singular bouncing effective field model with a non-standard kinetic term at both the background and perturbative levels. We find that these two settings share many important properties, including the result that they both generate scale-invariant scalar perturbations. This shows that some quantum gravity effects of the very early universe may be mimicked by effective field models. ",
        "introduction": "\\label{s.intro}  High precision observations of the cosmic microwave background (CMB) \\cite{Spergel:2003cb, Ade:2013zuv} and the large scale structure provide strong evidence that the structure of our universe was seeded by a nearly scale-invariant power spectrum of primordial curvature perturbations. Such a primordial power spectrum can be causally generated from quantum vacuum fluctuations during an era of cosmic inflation \\cite{Mukhanov:1981xt} where it is suggested that our universe experienced an exponentially accelerating expansion right after the initial big-bang singularity \\cite{Guth:1980zm}. Alternatively, as suggested in \\cite{Wands:1998yp, Finelli:2001sr}, a scale-invariant power spectrum of primordial perturbations can also be achieved during a matter-dominated contracting epoch, with the current expanding universe following a bouncing phase. A model of this type with an initial matter-dominated contraction and a non-singular bounce is called the matter bounce scenario, and provides an alternative to inflation for generating the observed spectrum of primordial fluctuations (see \\cite{RHBrevs} for review articles on this topic).  The bouncing phase that connects the contracting and expanding branches can in principle be either singular (as in the original ekpyrotic scenario \\cite{Khoury:2001wf}) or non-singular. Non-singular bouncing cosmologies appear in several settings where either the gravitational sector is modified as in Ho\\u{r}ava gravity \\cite{Calcagni:2009ar, Brandenberger:2009yt}, torsion gravity \\cite{Cai:2011tc, Poplawski:2011jz} and non-local gravity \\cite{Biswas:2005qr}, or the matter field violates positive energy conditions like in the quintom bounce \\cite{Cai:2007qw, Cai:2007zv}, the ghost condensate bounce \\cite{Buchbinder:2007ad, Lin:2010pf} and the Galileon bounce \\cite{Qiu:2011cy} models.  For a recent review of non-singular bouncing cosmologies, see \\cite{Novello:2008ra}.  Non-singular bounces also generically appear in loop quantum cosmology (LQC) \\cite{Singh:2009mz}, where the variables and quantization techniques of loop quantum gravity are used to study quantum gravity effects in cosmological space-times \\cite{Bojowald:2006da, Ashtekar:2011ni}. In LQC, the classical big-bang singularity  is replaced by a quantum bounce that occurs when the space-time curvature reaches the Planck scale \\cite{Ashtekar:2006wn}. Moreover, the full quantum dynamics of semi-classical states are approximated to a high degree of accuracy by a simple set of effective equations \\cite{Ashtekar:2006wn, Taveras:2008ke, Rovelli:2013zaa}. LQC is compatible with many interesting cosmological scenarios, including inflation \\cite{Ashtekar:2009mm}, the matter bounce \\cite{WilsonEwing:2012pu}, as well as the ekpyrotic scenario \\cite{Wilson-Ewing:2013bla}, and therefore it is interesting to compare the predictions coming from the effective field models described above and from LQC.  As we shall see, many of the predictions are very similar.  An important conceptual issue for bouncing cosmologies is the well-known Belinsky-Khalatnikov-Lifshitz (BKL) anisotropic instability \\cite{BKL}. The BKL instability appears in contracting cosmologies as the effective energy density contributed by the back-reaction of anisotropies increases faster than the energy densities of the dust and radiation matter fields. Therefore, in order to have a bounce that is approximately isotropic, it is necessary to fine-tune the initial conditions to be nearly perfectly isotropic in order to ensure that anisotropies never dominate. However, this problem is avoided in the ekpyrotic scenario where a scalar field with a steep and negative-valued potential always dominates over anisotropies in a contracting universe \\cite{Erickson:2003zm} and so it is justified to neglect anisotropies in the presence of an ekpyrotic scalar field.  Recently, it has been shown how it is possible to combine an era of ekpyrotic contraction with a non-singular bounce by introducing a scalar field with a Horndeski-type non-standard kinetic term and a negative exponential potential \\cite{Cai:2012va}. Furthermore, one may include a regular dust field and assume the universe began in a state of matter-dominated contraction thus combining the matter bounce with the ekpyrotic scenario. It can be explicitly checked that anisotropies remain small throughout the entire cosmological evolution of the matter-ekpyrotic bounce model, including at the bounce point \\cite{Cai:2013vm}, and therefore this model successfully avoids the BKL instability that arises for a large family of non-singular bounce models, as pointed out in \\cite{Xue:2010ux}. Among many possible implementations of the matter bounce, a concrete realization of the matter-ekpyrotic bounce was constructed in \\cite{Cai:2013kja} that involves two matter fields with one being the scalar field that causes the bounce and the other representing the matter field that is dominant at the beginning of the contracting phase. Note that this effective field theory model of a non-singular bounce can also be developed into a super-symmetric version \\cite{Koehn:2013upa}.  In this paper, we shall study the matter-ekpyrotic bounce within the context of LQC and compare its dynamics to those found in the effective field approach developed in \\cite{Cai:2012va}. We find that the evolution of both the background and perturbative degrees of freedom are comparable in these two settings, and that in particular the dynamics of the primordial cosmological fluctuations in these two frameworks possess very similar features. For example, the sound speed squared, which characterizes the propagations of these fluctuation modes, becomes negative during the bouncing phase in both settings. This indicates that what appeared to be an instability in the effective field theory may in fact mimic a quantum gravity effect. In addition, the scale-dependence of the power spectrum studied in the effective field approach coincides with the one obtained in LQC. These similarities indicate that some of the quantum gravity effects of the very early universe can indeed be reproduced by effective field theory models.  In addition, since the ekpyrotic scalar field dominates the dynamics as the bounce is approached, the BKL instability is avoided as the back-reaction due to anisotropies will remain negligible compared to the energy density of the ekpyrotic scalar field. Because of this property of the matter-ekpyrotic bounce model, it is possible to neglect anisotropies and take the background to be the simplest homogeneous cosmology, namely the isotropic flat Friedmann-Lema\\^itre-Robertson-Walker (FLRW) universe where the dynamics, of both the background cosmology and the perturbations, are significantly simpler.  The paper is organized as follows. In Sec.\\ \\ref{s.eff} we briefly review the matter-ekpyrotic bounce effective field theory model with the scalar field with a non-standard Horndeski-type kinetic term. Then in Sec.\\ \\ref{s.hom} we study the evolution of the matter-ekpyrotic homogeneous cosmology in LQC where the scalar field now has the standard kinetic term (and the potential is taken to be the same). We show explicitly how a non-singular bounce occurs near the Planck regime where quantum gravity effects become important. Afterwards, in Sec.\\ \\ref{s.scal} we analyze the dynamics of cosmological perturbations both analytically and numerically in the LQC matter-ekpyrotic model by tracking each Fourier mode throughout the cosmic evolution and show how the modes that exit the Hubble radius during the initial period of matter-dominated contraction are scale-independent. We close with some concluding remarks in Sec.\\ \\ref{s.disc}.  In this paper, we use units where $c = \\hbar = 1$ but leave the reduced Planck mass $M_{\\rm Pl}  = 1/\\sqrt{8\\pi G_N}$ (with $G_N$ being Newton's gravitational constant) explicit.  Also, the normalization of the scale factor is chosen so that its value at the bounce time $t=0$ is $a_{B} \\equiv a(t=0) = 1$. ",
        "conclusions": "\\label{s.disc}  In LQC, the big-bang and big-crunch singularities are generically resolved without any need to introduce the non-standard kinetic operators that are required in the effective field approach. In this paper we study the cosmological dynamics of a canonical scalar field with an ekpyrotic-like potential and a dust field in the context of LQC. With this choice of matter content, the universe initially starts its evolution with an epoch of matter-dominated contraction, then it transits into a period of ekpyrotic contraction, and afterwards experiences a smooth bouncing phase which connects to a regular thermal expansion. The background evolution obtained based on quantum gravity effects in LQC is very similar to that described by the effective field approach involving non-canonical kinetic operators.  One promising feature of the matter-ekpyrotic bounce cosmology is that it generates a scale-invariant power spectrum of primordial curvature perturbations. The scale-invariant perturbations are generated in the matter-dominated contracting phase and their amplitude is determined by the Hubble rate at the end of this phase, which in our model is denoted by $H_e$. We show that this scale invariance is well preserved across the bounce and compare this with the results found in the effective field approach (for example see the study of \\cite{Finelli:2001sr} based on matching conditions \\cite{Hwang:1991an, Deruelle:1995kd} and \\cite{Cai:2007zv, Cai:2008qw, Xue:2013bva} using various field theory models). Interestingly, the effective theory model captures much of the relevant physics of LQC even though the effects of LQC only become important in the deep Planck regime.  This is somewhat surprising as effective field theories are typically only expected to be reliable when quantum gravity effects are small, but here we have seen that the main properties of the evolution through the bounce in LQC ---of both the homogeneous background and the perturbations--- are captured in the effective theory models, even though these effects occur in the deep non-perturbative regime of LQC.  In the effective theory approach, the bounce usually occurs well below the Planck regime in order to ensure that quantum gravity effects will not ruin the validity of the effective theory.  However, we have seen here that effective theory models in fact capture the key qualitative features of LQC, and therefore can act as toy models for full quantum gravity when the bounce energy scale of these models is in the Planck regime. This suggests that effective field theory models in cosmology may be more reliable at high energy scales than initially expected."
    },
    "1402/1402.2525_arXiv.txt": {
        "abstract": "{ The COMPTEL experiment on CGRO observed the \\gray\\ sky at energies from 0.75 MeV to 30 MeV between April 1991 and June 2000. COMPTEL detected many \\gray\\ sources, among them an unidentified one labeled GRO~J1823-12. It is located near l/b$=$17.5\\deg/-0.5\\deg\\ and positionally consistent with the prominent \\gray\\ binary LS~5039. } { LS~5039 was established as a \\gray\\ source at TeV energies by HESS and at GeV energies by Fermi/LAT during recent years, whose \\gray\\ radiation is modulated along its binary orbit. Given this new information we reanalysed the COMPTEL data of GRO~J1823-12 including an orbital resolved analysis. } { We applied the standard methods, proper event selections and data binning with subsequent maximum-likelihood deconvolution, to analyse the COMPTEL data. In addition we developed a tool to select and bin the COMPTEL data in a phase-resolved manner. We present the orbit-averaged as well as orbit-resolved MeV analyses, lightcurves and spectra, and put them into multifrequency context. } { The COMPTEL data show a significant MeV source, which is positionally consistent with LS~5039, however also with other closeby Fermi/LAT sources. The orbit-resolved analysis provides strong evidence, at about the 3\\sig\\ level, that the MeV flux of GRO~J1823-12 is modulated along the binary orbit of about 3.9 days of LS~5039. We show that at MeV energies the source is brighter at the orbital part around the inferior conjuction than at the part of the superior conjunction, being in phase with X-rays and TeV \\grays, however being in anti-phase with GeV \\grays. The high-energy SED (X-rays to TeV \\grays) shows the high-energy emission maximum of LS~5039 at MeV energies. } { We conclude that the COMPTEL source GRO~J1823-12 is the counterpart of the microquasar LS~5039, at least for the majority of its MeV emission. The COMPTEL fluxes, put into multifrequency perspective, provide new constraints on the modelling of the high-energy emission pattern of the \\gray\\ binary LS~5039. } ",
        "introduction": "The Compton Gamma-Ray Observatory (CGRO) observed the universe in \\gray\\ energies with unprecedented sensitivity for more than 9 years between its launch in April 1991 and its reentry into the earth atmosphere in June 2000. CGRO carried the four \\gray\\ experiments BATSE, OSSE, COMPTEL and EGRET, which provided many new and exciting results in the regime of the \\gray\\ band, accesible only from space. In particular, many point sources between \\sm 1~MeV and \\sm10~GeV were detected, of which surprisingly a big fraction remained unidentified until the end of the CGRO mission. For example, out of the 271 EGRET \\gray\\ sources listed in the 3rd EGRET source catalogue for energies above 100~MeV, 171 remained unidentified \\citep{Hartman99}.  The COMPTEL experiment was sensitive to soft \\grays\\ (0.75-30~MeV) and finally opened the MeV-band, basically unexplored before CGRO, as a new astronomical window. Apart from \\gray\\ bursts and AGN, the majority of the COMPTEL sources are unidentified objects. The first COMPTEL source catalogue lists 10 AGN and 9 unidentified \\gray\\ sources \\citep{Schoenfelder00}. One of these unidentified sources is GRO~J1823-12, which is significantly visible in time-averaged COMPTEL maps of all standard energy bands between 1 and 30 MeV \\citep[e.g.][] {Collmar00a}. GRO~J1823-12 is located near l/b$=$17.5\\deg/-0.5\\deg, spatially consistent with three EGRET \\gray\\ sources, the most prominent one being 3EG~J1824-151 which was proposed as the counterpart of the microquasar LS~5039 \\citep{Paredes00}. Results of  preliminary analyses of the COMPTEL data on GRO~J1823-12 between 1991 and 1997 were reported by \\citet{Strong01} and \\citet{Collmar03}. The source showed a hard power-law spectrum with photon index of \\sm1.6 in the energy band between 1-30~MeV. However, due to the large error location area of the COMPTEL source, the source remained unidentified.  LS~5039 \\citep{Stephenson71}, a luminous star of the southern milky way (l/b$=$16.88\\deg/-1.29\\deg), was first identified as a high-energy source by a cross-correlation with unidentified ROSAT X-ray sources \\citep{Motch97}. They suggested LS 5039 to be a high-mass X-ray binary system. The detection of LS~5039 by the High Energetic Spectroscopic System (HESS) at energies above 250 GeV \\citep{Aharonian05} proved LS~5039 being even a \\gray\\ source. The HESS measurements showed that the flux and the energy spectrum of LS~5039 is modulated with the orbital period of the binary system of \\sm3.9 days \\citep{Aharonian06a}. At present LS~5039 is well established as being a high-energy source, e.g. at energies above 100~MeV by Fermi/LAT \\citep{Abdo09, Hadasch12}, at hard X-rays (25 - 200~keV) by INTEGRAL \\citep{Hoffmann09} and at X-ray energies (1 - 10~keV) by e.g. {\\it Suzaku} \\citep{Takahashi09}. The microquasar is significantly detected at these energies, including the orbital modulation of its flux and energy spectrum.  Because this peculiar time variability of LS~5039 became known during last years and was detected in COMPTEL's neighboring energy bands, we reanalysed our COMPTEL data on GRO~J1823-12. The detection of such a time signature in the COMPTEL MeV data would (at least partly) identify this COMPTEL source as counterpart of LS~5039. In this paper we report the analysis results and discuss them.  The paper is organized as follows: in Sect. 2 we briefly describe the COMPTEL instrument and the applied data analysis methods, the analysed observations are summarized in Sect. 3. In Sect. 4 we present the analysis results and a high-energy SED in Sect.~5. In Sect. 6 we discuss our results and finally summarise and conclude in Sect.~7. ",
        "conclusions": "Summary and conclusion}   LS~5039 is now an established \\gray\\ source at energies above 100 MeV, for which crucial spectral and timing infos became available in recent years. Because LS~5039 is spatially coincident with the known but unidentfied COMPTEL source GRO~J1823-12, we reanalysed the COMPTEL data of this sky region by assuming that LS~5039 is the counterpart of GRO J1823-12. We report the data analysis work in this paper.  We provide strong evidence for LS~5039 being the counterpart of GRO~J1823-12. Individual statistical tests are not totally convincing by reaching individually up to 3\\sig\\ only. However from the sum of the analyses, we conclude that LS~5039 is -- at least for the major fraction -- the counterpart of the MeV source. The derived absolute flux values, fitted at the sky position of LS~5039 and assumed to be from LS~5039, may slightly be too high because we cannot exclude underlying emission from further \\gray\\ sources which COMPTEL is unable to resolve. From the orbital lightcurve, we estimate them to be up to one third at most by assuming that the minimum emission is completely due to other sources. For the relative fluxes we show that the source is brighter during {\\it INFC} than during {\\it SUPC}, which shows the dominance in radiation for LS~5039 has to flip back from {\\it SUPC} at energies above 100~MeV to {\\it INFC} between 30 and 100~MeV. This is an important result which provides a significant constraint for the source modelling. Adding the COMPTEL fluxes to the high-energy SED shows that the emission maximum of LS~5039 is at MeV energies. While the {\\it SUPC} emission almost smoothly turns over from the Fermi/LAT band to the COMPTEL band, the {\\it INFC} emission seems to show a strong break between the COMPTEL and the Fermi/LAT band. This behaviour provides another important constraint for modelling the \\gray\\ binary.  In summary, the COMPTEL data have enlighten a part of the yet ``dark'' region in the SED of the \\gray\\ binary LS~5039 thereby providing new insights in its emission processes by delivering new and important constraints for the source modelling."
    },
    "1402/1402.0520_arXiv.txt": {
        "abstract": "We discuss how dynamical dark energy universes with exotic singularities may be distinguished from the standard $\\Lambda$CDM model on the basis of their redshift drift signal, for which measurements both in the acceleration phase and in the deep matter era will be provided by forthcoming astrophysical facilities. Two specific classes of exotic singularity models are studied: sudden future singularity models and finite scale factor singularity models. In each class we identify the models which can mimic $\\Lambda$CDM and play the role of dark energy as well as models for which redshift drift signals are significantly different from $\\Lambda$CDM and the test can differentiate between them. ",
        "introduction": "\\setcounter{equation}{0}  Exotic singularities are all the non-big-bang/-big-crunch types of singularities which started being investigated after the discovery of the accelerated expansion of the universe \\cite{supernovaold}. Their classification was first attempted in Ref. \\cite{nojiri} and further extended in Refs. \\cite{AIP2010,sesto}. The list is pretty long and includes a big-rip \\cite{phantom}, a little rip \\cite{LRip}, a pseudo-rip \\cite{PRip}, a sudden future singularity \\cite{BGT,Varun,barrow04}, a generalized sudden future singularity \\cite{GSFS}, a finite scale factor singularity \\cite{nojiri,JCAP12,ABTD}, a big-separation \\cite{nojiri}, a $w$-singularity \\cite{wsing}, and their generalizations \\cite{yurov}. All of them but big-rip and little-rip are geodesically complete and do not fulfil the standard requirement for being a singularity required by the Hawking-Penrose theorems. The geodesic equations do not feel these singularities \\cite{LFJ} and due to this, even extended objects may pass through them without being destroyed \\cite{adam}. This property allows some of exotic singularities to show up in the near future with no contradiction to observational data and in such a way they may serve as a kind of dark energy which leads to the acceleration of the universe.  In this paper two particular examples of exotic singularities will be studied. The first is a sudden future singularity (SFS, which seemingly appear uniquely in loop quantum cosmology \\cite{LQC}) which has been shown to be observationally allowed \\cite{DHD,GHDD,DDGH,laszlo}, and the second is a finite scale factor singularity (FSFS) which was also checked against standard cosmological data (supernovae, baryon acoustic oscillations, CMB shift parameter) \\cite{JCAP12,ABTD}. Some other before mentioned types of exotic singularities have also been tested (e.g. \\cite{f(rho)}). Here we are going to concentrate on an astrophysical observable known as the redshift drift \\cite{sandage}. This test has been discussed in many other cosmological frameworks \\cite{Quercellini12,pauline12}, but not yet in the context of exotic singularities. We want to fill in this gap and discuss it in the present contribution.  The paper is organized as follows. In Section \\ref{exotic} we briefly review the basic properties of exotic singularities under study. In Section \\ref{drift} we provide a short overview of redshift measurements, and in Sect. \\ref{results} we discuss its signatures for these singularities. In Section \\ref{rac} we present our conclusions. ",
        "conclusions": "\\label{rac} \\setcounter{equation}{0}  We have presented some viable examples of exotic singularity universes and studied the redshift drift effect which allows to look at two different past light cones of an observer placed on Earth separated by the time of 20 years. Two particular models have been investigated: the ones which admit sudden future (or pressure) singularities and another ones which admit finite scale factor (of energy density and pressure) singularities.  Our analysis shows that the range of redshift signal is made discrepant from $\\Lambda$CDM is model-dependent: some models in this class are easier to distinguish at low redshifts (within reach od SKA or CHIME) while for others to high redshifts probed by ELT-HIRES will be ideal.  We have shown that both these models, for special values of the parameters (chosen for models marked as SFS3 and FSFS1), can mimic the redshift drift behavior of $\\Lambda$CDM models. In other words, the matter which is responsible for the singularities behaves as dark energy. As for other values of the parameters (models SFS1, FSFS3) the redshift drift test can play the role of a differential test between these models and $\\Lambda$CDM, though leaving them still good candidates for dark energy."
    },
    "1402/1402.0466_arXiv.txt": {
        "abstract": "{We believe that the Dark Matter (DM) search should be expanded into the domain of detectors sensitive to molecular excitations, and so that we should create detectors which are more sensitive to collisions with very light WIMPs. In this paper we investigate in detail di-atomic molecules, such as Fused Silica material with large OH-molecule content, and water molecules. Presently, we do not have suitable low cost IR detectors to observe single photons, however some OH-molecular excitations extend to visible and UV wavelengths, and can be measured by Bialkali photocathodes. There are many other chemical substances with di-atomic molecules, or more complex oil molecules, which could be investigated also. This idea invites searches in experiments having large target volumes of such materials coupled to a large array of single-photon detectors with Bialkali or infrared-sensitive photocathodes.} ",
        "introduction": "\\begin{figure}[tbp] \\includegraphics[width=0.45\\textwidth]{Energy_band_structure_of_inorganic_scintillator.pdf} \\caption{Energy levels for an activated NaI(Tl) crystal, where Tl is the so-called activator. The Tl-levels lie somewhat below the conduction band so that the radiated photons are emitted with less energy than the band gap of the pure crystal.} \\label{fig:NaI_energy_levels} \\end{figure}  \\begin{figure}[tbp] \\includegraphics[width=0.4\\textwidth]{Vibration_spectra_of_diatomic_molecules.pdf} \\caption{Energy levels of diatomic molecules for a harmonic and an anharmonic oscillator, where the latter follows the more complex Morse potential; $\\nu$ is the vibrational quantum number ($\\nu$ = 0,1,2,...). The $\\nu$ = 0 level is the vibrational ground state. Unlike the energy levels of the harmonic oscillator potential, which are evenly spaced by $\\hbar\\omega$, the Morse potential level spacing decreases as the energy approaches the dissociation energy $D_{e}$ [\\cite{Morse_2014}].} \\label{fig:Vibration_spectra_of_diatomic_molecules} \\end{figure}  A recent paper [\\cite{Vavra_2014_a}] discussed the possibility that the DAMA experiment [\\cite{Bernabei_2013}] observes the effect of a light WIMP striking a hydrogen atom and creating a proton recoil, which then excites the NaI(Tl) crystal in which it is located. The target hydrogen atom is present as OH or H$_2$O contamination. The \"classical\" scintillation mechanism for atomic level excitation in NaI(Tl) is shown schematically in Fig.~\\ref{fig:NaI_energy_levels}. The primary energy deposit has to exceed the band gap of NaI, which is ~3.9 eV; thallium is added to lower the emission energy.  The present paper considers another, even more sensitive, method to detect Dark Matter via molecular excitations (a preliminary version was presented in Ref.~[\\cite{Vavra_2014_b}]). A very small energy-deposit, at a level of 2-3~eV, is needed to excite such vibrations; this is a factor of 2 lower in energy than the classical scintillation mechanism. Many di-atomic molecules will vibrate when hit by a photon (or neutron or a WIMP) and some small fraction of such single photon excitations can be detected by a PMT working in the Bialkali regime. However, most of the emitted energy is in the IR wavelength region, or is in the form of heat; the emission of visible photons is suppressed by at least 4-6 orders of magnitude compared to IR photons. Because the efficiency for Dark Matter to generate visible photons in water or ice is expected to be very small, a large detector volume is required, together with an extensive PMT coverage with single photon sensitivity. Examples of possible experiments to do this measurement are the Ice Cube, BaBar DIRC and the Super-Kamiokande experiment.  \\begin{table} \\caption{A simple calculation of the transition wavelength for several frequency overtones of the OH-radicals. The last two modes correspond to visible wavelengths. Higher modes can reach the UV regime [\\cite{Chemistry_lectures}].} \\label{table:1}                     % \\resizebox{9cm}{!} { \\begin{tabular}{c c c}                  % \\hline\\hline                            % OH-band identity  & Transition &  Calculated wavelength [nm]  \\\\     % \\hline                                  % $\\nu_{1}$ & 0 $\\rightarrow$ 1 & 2803 \\\\ 2$\\nu_{1}$ & 0 $\\rightarrow$ 2 & 1436 \\\\ 3$\\nu_{1}$ & 0 $\\rightarrow$ 3 & 980 \\\\ 4$\\nu_{1}$ & 0 $\\rightarrow$ 4 & 755 \\\\ 5$\\nu_{1}$ & 0 $\\rightarrow$ 5 & 619.5 \\\\ \\hline                                   % \\end{tabular} } \\end{table}  Di-atomic molecules are ideal for such studies, since they are simpler to understand, and have been studied extensively by chemists [\\cite{Morse_1929}]. The energy levels of di-atomic molecules can be calculated using the Schroedinger equation description of a harmonic oscillator, which has been described in many textbooks - $e.g.$ see Refs. [\\cite{Morse_2014}] and [\\cite{Chemistry_lectures}]. Figure~\\ref{fig:Vibration_spectra_of_diatomic_molecules} and Table~\\ref{table:1} provide examples of energy levels in a typical di-atomic molecule, for a perfect harmonic oscillator, and for an anharmonic oscillator following the Morse potential [\\cite{Morse_2014}]. The most significant difference between the two potentials is a transition selection rule. The harmonic oscillator allows only transitions obeying $\\bigtriangleup\\nu=\\pm1$. In contrast, the anharmonic oscillator allows the transitions $\\bigtriangleup\\nu=\\pm1$, $\\pm2$, $\\pm3$, etc. These transitions, $\\nu_{1}$, 2$\\nu_{1}$, 3$\\nu_{1}$, etc., are called overtones. We see that although the fundamental mode corresponds to an IR wavelength, the higher overtones correspond to visible or even UV wavelengths. This can provide a path to detection using Bialkali photocathode-based PMTs; single-photon IR-sensitive detectors are more difficult to utilize at present for this purpose.  \\begin{figure}[tbp] \\includegraphics[width=0.4\\textwidth]{OH_energy_levels.pdf} \\caption{A schematic picture of WIMP-scattering excitation of energy levels in the OH-molecule, with subsequent de-excitation via fluorescent photons at either $\\sim$282~nm or $\\sim$310~nm.} \\label{fig:OH_levels} \\end{figure}  We note that the OH-molecule has been studied extensively by many chemists using laser-induced fluorescence. An example of such a fluorescence measurement is the OH-molecule excitation by a 282~nm dye laser, followed by the observation of the 310~nm wavelength with a PMT with a notch filter [\\cite{Smith_1990}]. This method was used to determine traces of OH-radicals in the atmosphere [\\cite{Matsumi_2002}]. Figure~\\ref{fig:OH_levels} shows schematically how WIMP scattering could excite energy levels in the OH-molecule, with a subsequent de-excitation via fluorescent photons at either $\\sim$282~nm or $\\sim$310~nm.  However, the picture outlined above is too simplistic, since the actual energy levels can be affected by the substance in which the OH-radical is located. In this paper we discuss the example of Fused Silica loaded with OH-molecules. Fused Silica is a good example because it has been investigated previously in great detail. It is an ideal material in which to study the effects of the OH-molecular presence, because the basic material is very pure. In addition, Fused Silica can be prepared in the form of long fibers, which allows the accurate measurement of absorption resonances by using a monochromator. The next chapter will deal with this material in more detail.  Although Fused Silica is a well-understood material, pure water is even better understood. We have looked into its molecular vibration response also. ",
        "conclusions": "This paper provides a scheme within which molecular vibrations might play a role in the discovery of a light WIMP. We have described in detail the vibrations of OH-molecules in Fused Silica and in pure water. A very small energy deposit at the level of 2-3~eV is needed to excite such vibrations, with the corresponding de-excitations detectable by a Bialkali photocathode. This is a factor of 2 lower in energy than is needed to excite the classical scintillation mechanism in a NaI(Tl) crystal. Since the efficiency of the light emission via this process is small, a very large radiator volume of $e.g.$ water, Fused Silica or oil is required, together with extensive detector coverage with single-photon detection capability."
    },
    "1402/1402.5075_arXiv.txt": {
        "abstract": "We discuss how Starobinsky-like inflation may emerge from dilaton dynamics in brane cosmology scenarios based on string theory, in which our universe is represented as a three-brane. The effective potential may acquire a constant term from a density of effectively point-like non-pertubative defects on the brane. Higher-genus corrections generate corrections to the effective potential that are exponentially damped at large field values, as in the Starobinsky model, but at a faster rate, leading to a smaller prediction for the tensor-to scalar perturbation ratio $r$. This may be compensated partially by logarithmic deformations on the world-sheet due to recoil of the defects due to scattering by string matter on the brane, which tend to enhance the tensor-to-scalar ratio.  Quantum fluctuations of the ensemble of D-brane defects during the inflationary period may also enhance the tensor-to-scalar ratio. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.2655_arXiv.txt": {
        "abstract": "We use high-quality, medium-resolution {\\it Hubble Space Telescope}/Cosmic Origins Spectrograph (\\HST/COS) observations of 82 UV-bright AGN at redshifts $z_{\\rm AGN}<0.85$ to construct the largest survey of the low-redshift intergalactic medium (IGM) to date: 5138 individual extragalactic absorption lines in \\HI\\ and 25 different metal-ion species grouped  into 2611 distinct redshift systems at $z_{\\rm abs}<0.75$ covering total redshift pathlengths $\\Delta z_{\\rm HI}=21.7$ and $\\Delta z_{\\rm OVI}=14.5$.  Our semi-automated line-finding and measurement technique renders the catalog as objectively-defined as possible.  The cumulative column-density distribution of \\HI\\ systems can be parametrized $d{\\cal N}(>N)/dz=C_{14}(N/10^{14}$~\\cd$)^{-(\\beta-1)}$, with $C_{14}=25\\pm1$ and $\\beta=1.65\\pm0.02$.  This distribution is seen to evolve both in amplitude, $C_{14}\\propto(1+z)^{2.3\\pm0.1}$, and slope $\\beta(z)=1.75-0.31\\,z$ for $z\\le0.47$.  We observe metal lines in 418 systems, and find that the fraction of IGM absorbers detected in metals is strongly dependent on \\NHI.  The distribution of \\OVI\\ absorbers appear to evolve in the same sense as the \\Lya\\ forest.  We calculate contributions to $\\Omega_b$ from different components of the low-$z$ IGM and determine the \\Lya\\ decrement as a function of redshift.  IGM absorbers are analyzed via a two-point correlation function in velocity space.  We find substantial clustering of \\HI\\ absorbers on scales of $\\Delta v=50-300$ \\kms\\ with no significant clustering at $\\Delta v\\ga1000$~\\kms.  Splitting the sample into strong and weak absorbers, we see that most of the clustering occurs in strong, $N_{\\rm HI}\\ga10^{13.5}$ \\cd, metal-bearing IGM systems. The full catalog of absorption lines and fully-reduced spectra is available via the Mikulski Archive for Space Telescopes (MAST) as a high-level science product at {\\tt http://archive.stsci.edu/prepds/igm/}. ",
        "introduction": "The low-redshift intergalactic medium (IGM) holds many important clues to complete our understanding of cosmology.  Even after nearly 14 Gyr of evolution, only a small fraction of baryonic matter has collapsed into luminous objects (galaxies, groups, clusters) while $\\sim80\\%$ or more still exists as a diffuse, often unvirialized IGM, the distribution and characteristics of which are only beginning to be measured \\citep[e.g.,][]{Shull12a}.  Furthermore, there is a complicated and poorly understood interplay between the IGM, circumgalactic medium (CGM), and stars and gas in galaxies.  These gaseous reservoirs provide raw material which is subsequently formed into stars and galaxies.  These, in turn, enrich the IGM via outflows driven by supernovae, radiation pressure, and active galactic nuclei \\citep[AGN][]{OppenheimerDave08,Smith11}.  Diffuse intergalactic gas is currently quite difficult to observe in emission \\citep[Frank \\etal\\ 2012; but see also][]{Steidel11,Martin14a,Martin14b}.  The most sensitive method for detecting most of the gas is through absorption-line spectroscopy using bright background objects (typically AGN) to provide an ultraviolet continuum.  The highest concentration of strong gas-diagnostic lines is in the rest-frame far-ultraviolet (FUV) band from $\\sim2000$~\\AA\\ shortward to the Lyman edge at 912~\\AA.  Investigating the FUV at low redshift requires optimized spectrographs above Earth's UV-blocking atmosphere.  Thus, there have been a series of space-based UV spectrographs, both as primary-science instruments on space-borne observatories ({\\it Copernicus, International Ultraviolet Explorer, Hopkins Ultraviolet Telescope, Far-Ultraviolet Spectroscopic Explorer}) and instruments installed aboard the {\\it Hubble Space Telescope} or \\HST: Faint Object Spectrograph (FOS), Goddard High-Resolution Spectrograph (GHRS), Space Telescope Imaging Spectrograph (STIS), and now the Cosmic Origins Spectrograph (COS).  COS is the fourth-generation UV spectrograph onboard \\HST\\ and is optimized for medium-resolution ($\\lambda/\\Delta\\lambda\\approx18,000$, $\\Delta v\\approx17$ \\kms) spectroscopy of point sources in the 1135--1800~\\AA\\ band \\citep{Green12,Osterman11}.  COS has an effective area that is an order of magnitude larger in the FUV ($\\lambda<1800$~\\AA) than previous spectrographs.  Furthermore, the excellent scattered light control and low background detectors of COS mean that fainter objects ($F_\\lambda\\la10^{-14}$ \\flux) can be observed than with previous instruments, and higher quality spectra of bright targets can be obtained in much less time.  In its first five years of science operations \\citep{Green12}, COS accumulated an unprecedented archive of hundreds of AGN spectra collected under a broad range of scientific programs from AGN physics to studies of the interstellar medium (ISM) in our own Galaxy.  Regardless of their original scientific intent, many of these spectra are suitable for quasar absorption-line studies of the IGM.  In this paper, we build on the heritage of many previous low-$z$ IGM absorber catalogs from \\HST/FOS \\citep{Bahcall93,Bahcall96,Jannuzi98,Weymann98}, \\HST/GHRS \\citep{Penton1}, \\FUSE\\ \\citep{Danforth05,Danforth06}, and \\HST/STIS \\citep{Penton04,Lehner07,DS08,Tripp08,Danforth10a,Tilton12}.  This current survey represents the largest sample of low-$z$ absorbers to date and is more sensitive than previous studies in most cases.  It was designed to be a general-purpose survey with applications to a wide variety of extragalactic astrophysics: 82 extragalactic sight lines covering a combined redshift pathlength in \\HI\\ of $\\Delta z=21.7$, 5138 intergalactic absorption lines comprising 2611 distinct redshift systems, and detections of 25 different metal-ion species.  We describe the sample selection criteria, data reduction methods, and semi-automated measurement techniques in Section~2.  We present the catalog and substantial electronic resources available to users of the survey in Section~3.  Overall results from the survey are given in Section~4.  In Section~5, we present some of the more detailed findings from the survey, including the evolution of \\HI\\ and metals in the IGM at $z\\le0.47$ and the radial-velocity clustering properties of IGM absorbers (the two-point correlation function).  Important survey parameters, fitted quantities, and initial findings are summarized in Section~6. ",
        "conclusions": ""
    },
    "1402/1402.3245_arXiv.txt": {
        "abstract": "{} {The presence of a magnetic field can have a strong impact on the evolution of a binary star. However, only a dozen of magnetic OB binaries are known as of today and available to study this effect, including very few magnetic pulsating spectroscopic OB binaries. We aim at checking for the presence of a magnetic field in the B5IV hierarchical triple system \\object{HD\\,1976} with spectropolarimetric data obtained with Narval at the Bernard Lyot Telescope (TBL).} {We use orbital parameters of \\object{HD\\,1976} available in the literature to disentangle the Narval intensity spectra. We compute Stokes V profiles with the Least Square Deconvolution (LSD) technique to search for magnetic signatures. We then derive an estimate of the longitudinal magnetic field strength for each observation and for various line lists.} {Our disentangling of the intensity spectra shows that \\object{HD\\,1976} is a double-lined spectroscopic (SB2) binary, with the lines of the secondary component about twice broader than the ones of the primary component. We do not identify the third component. Moreover, we find that clear magnetic signatures are present in the spectropolarimetric measurements of HD\\,1976 and seem to be associated with the primary component. We conclude that \\object{HD\\,1976} is a magnetic slowly-pulsating double-lined spectroscopic binary star, with an undetected third component. It is the second such example known (with \\object{HD\\,25558}).} {} ",
        "introduction": "The magnetic field of OB binaries has not been much explored so far, from either an observational or theoretical point-of-view, while the presence of a magnetic field in a binary could have an important impact on the evolution of the binary.  In low-mass stars, tidal interactions are expected to induce large-scale 3D shear and/or helical flows in stellar interiors that can significantly perturb their magnetic dynamo. Similar flows may also influence the fossil magnetic fields of high-mass stars. Moreover, in combination with tides, magnetically driven winds/outflows in binary systems have long been suspected to be responsible for their orbital evolution, while magnetospheric interactions have been proposed to enhance stellar activity. In addition, the incidence of magnetic stars among OB binary systems provides a basic constraint on the detailed fossil origin of the magnetic field and whether such magnetic fields suppress binary formation. However, the crucial observational constraints required to study these various issues are, at present, nearly nonexistent.  In the frame of the MiMeS (Magnetism in Massive Stars) project \\citep[e.g.][]{neiner2011,wade2013}, we checked for the presence of a magnetic field in HD\\,1976.  \\object{HD\\,1976} (V746\\,Cas) is a bright (V=5.6), slowly pulsating B (SPB) star \\citep{decat2007,dubath2011} of spectral type B5IV. Its main period of variation is 1.065 d \\citep[Hipparcos,][]{perryman1997}. \\cite{decat2007} and \\cite{mathias2001} found at least one additional period at 2.5 d and proposed  additional candidate periods. \\object{HD\\,1976} seems to be a young star of $\\sim$60 Myr \\citep{decat2007,tetzlaff2011}, but note that photometric age determination can lead to uncertainties for multiple systems.  \\object{HD\\,1976} is known to be a hierarchical triple system consisting of a visual binary with a magnitude difference of about 0.9 magnitude \\citep{perryman1997, mason2009}, one of the component being itself a SB1 binary \\citep{decat2007}. The angular separation of the visual components is 0.157 arcsec according to Hipparcos \\citep{perryman1997} and 0.111 arcsec according to \\cite{mason2009}. The period of the visual binary is about 150 years \\citep{abt2005}. The SB9 catalogue \\citep{pourbaix2004} provides a period for the SB1 between 25.4 and 27.8 d with an eccentricity between 0.12 and 0.2.  In this paper we report on Narval spectropolarimetric observations of HD\\,1976 (Sect.~\\ref{obs}). We use orbital parameters available in the literature to disentangle the spectra of the two components (Sect.~\\ref{binary}). We then analyse the magnetic measurements with the Least Square Deconvolution (LSD) technique (Sect.~\\ref{analysis}) and discuss the results (Sect.~\\ref{discuss}). ",
        "conclusions": "\\object{HD\\,1976} is a very interesting object as known magnetic pulsating spectroscopic binary are rather rare. The only other known SPB example is \\object{HD\\,25558} \\citep{sodor2014}.  We plan on acquiring new spectra of \\object{HD\\,1976} to derive a more precise orbital solution. We will then be able to disentangle the spectra more reliably and measure the magnetic field of the individual components. This is necessary to check whether the secondary component is also magnetic as well as to derive more accurate measurements of the magnetic field of the primary. Moreover, we will be able to test whether the magnetic component is part of the SB1 and is the pulsating component.  Finally, the BinaMIcS (Binarity and Magnetic Interactions in various classes of stars) project \\citep[see][]{neiner2013}, which relies on two large programs of observations with ESPaDOnS at CFHT and Narval at TBL, will allow us to acquire spectropolarimetric observations of many magnetic binary systems of all spectral types. These new observations will provide additional examples of magnetic massive binary stars, in particular systems with short orbital periods. Comparing HD\\,1976 with other magnetic massive binaries will be very useful to derive information about the role of magnetic field on binarity and vice-versa."
    },
    "1402/1402.2211.txt": {
        "abstract": "The mobility of O atoms at very low temperatures is not generally taken into account, despite O diffusion would add to a series of processes leading to the observed rich molecular diversity in space. We present a study of the mobility and reactivity of O atoms on an amorphous silicate surface. Our results are in the form of RAIRS and temperature-programmed desorption spectra of O$_2$ and O$_3$ produced via two pathways: O + O and O$_2$ + O, investigated in a submonolayer regime and in the range of temperature between 6.5 and 30 K. All the experiments show that ozone is formed efficiently on silicate at any surface temperature between 6.5 and 30 K. The derived upper limit for the activation barriers of O + O and O$_2$ + O reactions is $\\sim$ 150 K/k$_b$. Ozone formation at low temperatures indicates that fast diffusion of O atoms is at play even at 6.5 K. Through a series of rate equations included in our model, we also address the reaction mechanisms and show that neither the Eley Rideal nor the Hot atom mechanisms alone can explain the experimental values. The rate of diffusion of O atoms, based on modeling results, is much higher than the one generally expected, and the diffusive process proceeds via the Langmuir-Hinshelwood mechanism enhanced by tunnelling. In fact, quantum effects turn out to be a key factor that cannot be neglected in our simulations. Astrophysically, efficient O$_3$ formation on interstellar dust grains would imply the presence of huge reservoirs of oxygen atoms. Since O$_3$ is a reservoir of elementary oxygen, and also of OH via its hydrogenation, it could explain the observed concomitance of CO$_2$ and H$_2$O in the ices. ",
        "introduction": "The chemical processes taking place in interstellar clouds can be considered the origin of the molecular diversity in the Universe. A wealth of infrared, millimeter- and microwave-wavelength observations has provided evidence of rich molecular abundances within interstellar clouds. This observational evidence for the general interstellar medium (ISM), however, cannot be met by the known gas-phase reactions alone. That is why surface reactions are necessarily invoked for the formation of a growing number of molecular species. Atoms and molecules from the gas phase accrete and gather on the cold surfaces of interstellar dust grains, and eventually react after surface diffusion. In fact, some of the most abundant molecules in the Universe (such as H$_2$, H$_2$O or CO$_2$) are formed on dust grains.\\cite{Ti13} Particularly, hydrogenation of interstellar ices is known to induce the formation of species in the solid phase and, recently, O-atom additions as well were invoked for processes leading to an even richer molecular diversity.\\cite{Ti13} Due to the supposed high abundance and its certain reactivity, oxygen and its chemistry may then play a central role in astrochemistry.  Gas phase molecular oxygen has been detected\\cite{Li12, La07, Go11} in molecular clouds ($\\rho$ Ophiuchi A and  OMC-1) and astrochemical models of dark clouds predict that condensed oxygen is likely to be a major component of apolar ices. Nevertheless, the key O$_2$ molecule remains elusive in the ISM, very probably owing to its short lifetime,\\cite{Du12} and also because it can be easily consumed at the surface of dust grains by the two most abundant atomic species, H and O, forming H$_2$O and O$_3$ respectively, as main products. Water is the most abundant species in the solid phase, while solid ozone has not yet been observed in molecular clouds. To date, only a recent work\\cite{Ch01} presents spectra compatible with the presence of O$_3$ towards IC 5146. This apparent lack of ozone could be explained either by a detection bias as the broad silicate absorption band at 10 $\\mu$m can easily mask the 9.6 $\\mu$m ozone feature, or by an actual lack of ozone in the ice.  The non-detection of solid ozone in dense molecular clouds is consistent with its short lifetime on the surface of dust grains due to its high reactivity. Mokrane et al\\cite{Mo09} and Romanzin et al\\cite{Ro11} have shown that O$_3$ + H  is an efficient process under interstellar conditions and should be able to destroy most of the O$_3$ formed on the ice to produce water. Water formation on the grain surfaces occurs via hydrogenation of three different oxygen species:\\cite{La13,Ti82} O, O$_2$, and O$_3$. All these routes have been extensively studied in laboratory.\\cite{Io08, Mi08, Du10, Ch12}  Water is ubiquitous and its omnipresence is certainly a result of its stability: among the species made of the 2 dominant reactive atoms H and O, water is the most stable molecule. Figure~\\ref{fig:Enthalpy} represents all the stable molecules that are composed of O and H atoms. Their enthalpies of formation are represented vertically. We have drawn the different reactions identified experimentally, of which water seems to be the end of the chemical journey in the O and H world. However, to know how hydrogenation of oxygen takes place for forming water, it is necessary to understand how O$_2$ and O$_3$ are formed.   \\begin{figure} \\centering \\includegraphics[width=8.6cm]{Entalpia-eps-converted-to.pdf} \\caption{Schematic 3D representation of the water formation network via H-atom (grey circle) additions to O, O$_2$, and O$_3$ (red circles). Enthalpies of reactions grow along the vertical axis. Solid lines represent barrierless reactions, while dotted traces indicate reactions with an activation barrier.} \\label{fig:Enthalpy} \\end{figure} In 1930, S. Chapman discovered the mechanisms that produce the ozone layer in the Earth's stratosphere:\\cite{Ch68} UV photons striking oxygen molecules (O$_2$) split them into two oxygen atoms (O); atomic oxygen then combines with O$_2$ to create ozone. In turn, O$_3$ can be dissociated by UV light into a molecule of O$_2$ and an O atom, and so on in a continuing process called the ozone-oxygen cycle, creating an ozone layer in the stratosphere. Recently, ozone formation has been studied in the laboratory using supra thermal oxygen atoms generated by energetic electrons or ions.\\cite{Be05,Si07,En01} Jing et al\\cite{Ji12} also performed experiments on the formation of ozone on bare silicates, but our present work and analysis do not lead to the same conclusions. We attribute the detection of high temperature signals at mass 32 a.m.u. to the decomposition of O$_3$ inside the quadruple mass spectrometer, and not to the detection of O$_2$. In this paper we study the surface formation of O$_2$ and O$_3$ without the addition of energy, through the reactions:\\cite{Be94} \\begin{align} %XX\\begin{equation} O(^3P) + O(^3P) &+ M\\longrightarrow  O_2 +M  \\\\ O(^3P) + O_2(X^3 \\Sigma_{g}^-) &+ M \\longrightarrow   O_3(X^1 A_1) + M %XX\\end{equation} \\end{align}  The substrate, made of amorphous silicate,\\cite{No12} was held in the 5~--~30~K temperature range. Sub-monolayer conditions were used in all the experiments discussed below. This paper is organized as follows: the experimental set-up and methods are described in the next section. In Section 3, we present our results about O$_2$ and O$_3$ formation. In Section 4, we present a model that simulates our results and gives relevant energetic parameters. In the last Section, we discuss the main conclusions and astrophysical implications of this study. ",
        "conclusions": "In this paper we have shown that O$_3$ can be formed very efficiently via an atomic oxygen beam sent on silicate held at low temperatures (6-30 K). The reactions leading to ozone formation studied in this paper (O+O and O+O$_2$) appear to be barrierless; we have estimated an upper limit for the activation energies of reactions of 150 K/k$_b$. Conversely, the reaction O+O$_3$ has a high activation barrier -- lower limit $\\sim$ 2300 K/k$_b$ -- and it is not an efficient pathway for ozone destruction under our experimental conditions. In addition, the formation of ozone is favored by a very fast diffusion of oxygen atoms at low temperatures. The diffusive process of O atoms is likely to occur via quantum tunneling, as claimed in Minissale et al 2013,\\cite{Mi13} while the Hot Atom mechanism effects proved to be negligible.  From an astrophysical point of view, since the gas phase abundances of O and O$_2$ are elusive,\\cite{Je09, Wh10} it is difficult to put the O$_3$ formation in a simple interstellar contest. New and more detailed observational data are necessary to know the solid phase abundances of ozone. Anyway, we can fairly assume that in dense clouds, particularly UV-protected interstellar environments (A$_v$ $>$ 3), if comparable budgets of O atoms and H atoms are present,\\cite{Ca02} an O-addition chemistry competes with H additions and O$_3$ could be formed. The presence of ozone in the interstellar ices would confirm that the O$_3$+H pathway is the most important route leading to water in some interstellar environments. % (depend on O/H ratio). Not only ozone is important because it represents an efficient way to produce water, but also because it can be a reservoir of oxygen atoms; due to its low binding energy (1.1 eV), ozone can be easily processed by cosmic-rays and by the mean Galactic UV field,\\cite{Ti82, St91} producing O atoms that, for example, could react on the surface of dust grains with CO and produce CO$_2$.\\cite{Ma06,Ra11} Even without energetic events, CO+OH\\cite{Ob10, No11} and CO+O\\cite{Mi13b} reactions are believed to be sources of CO$_2$. The observed concomitance of CO$_2$ and H$_2$O in the ices\\cite{Wh10} can be more easily understood by assuming that this chemistry at the surface of dust grains is driven by the presence or absence of O$_3$. Ozone can be either an OH provider via its hydrogenation, or an O consumer upon its own formation (O $+$ O$_2$). Therefore, oxygen diffusion and reactivity on cold surfaces are also key factors to understanding the CO$_2$ and H$_2$O formation rates."
    },
    "1402/1402.3590_arXiv.txt": {
        "abstract": "We test the cosmological implications of studying galaxy clustering using a tomographic approach, by computing the galaxy two-point angular correlation function $\\omega(\\theta)$ in thin redshift shells using a spectroscopic-redshift galaxy survey. The advantages of this procedure are that it is not necessary to assume a fiducial cosmology in order to convert measured angular positions and redshifts into distances, and that it gives several (less accurate) measurements of the angular diameter distance $D_\\rmn{A}(z)$ instead of only one (more precise) measurement of the effective average distance $D_\\rmn{V}(z)$, which results in better constraints on the expansion history of the Universe. We test our model for $\\omega(\\theta)$ and its covariance matrix against a set of mock galaxy catalogues and show that this technique is able to extract unbiased cosmological constraints. Also, assuming the best-fit $\\Lambda$CDM cosmology from the cosmic microwave background measurements from the Planck satellite, we forecast the result of applying this tomographic approach to the final Baryon Oscillation Spectroscopic Survey catalogue in combination with Planck for three flat cosmological models, and compare them with the expected results of the isotropic baryon acoustic oscillation (BAO) measurements post-reconstruction on the same galaxy catalogue combined with Planck. While BAOs are more accurate for constraining cosmological parameters for the standard $\\Lambda$CDM model, the tomographic technique gives better results when we allow the dark energy equation of state $w_{DE}$ to deviate from $-1$, resulting in a performance similar to BAOs in the case of a constant value of $w_{DE}$, and a moderate improvement in the case of a time-dependent value of $w_{DE}$, increasing the value of the Figure-of-Merit in the $w_0-w_a$ plane up to $15\\%$. ",
        "introduction": "The study of large scale structure (LSS) has been of great importance for the advancement of our understanding of the Universe, characterising the distribution of structures, such as galaxies and voids, at large scales. Supported by the increasing amount of data from current and future large galaxy surveys, such as the Baryon Oscillation Spectroscopic Survey \\citep[BOSS; ][]{BOSS}, WiggleZ \\citep{WggZ}, the Dark Energy Survey \\citep[DES; ][]{DES}, the Hobby-Eberly Telescope Dark Energy Experiment \\citep[HETDEX; ][]{HETDEX}, the Large Synoptic Survey Telescope \\citep[LSST; ][]{LSST} and the EUCLID mission \\citep{EUCLID}, in combination with new and more precise measurements of the cosmic microwave background (CMB), the study of the LSS has a promising future in terms of shedding light on the nature of the Universe.  One of the most important cosmological probes of LSS is the signal of the baryon acoustic oscillations (BAO) measured in two-point statistics, such as the correlation function or the power spectrum. These oscillations occur because small primordial perturbations induce sound waves in the relativistic plasma of the early Universe \\citep{PeeblesYu70}, but later on at the recombination epoch ($z\\approx1000$), the sound speed suddenly decreases to the point that these waves stop propagating. Since the Universe has an appreciable fraction of baryons, cosmological theories predict that the BAO signal will also be imprinted onto the two-point statistics of the non-relativistic matter as an excess of clustering in the correlation function, or an oscillation in power in the power spectrum, at a given (fixed) scale, making it an ideal standard ruler.  In 1999, motivated by the results obtained from COBE of the primary temperature anisotropy in the CMB \\citep{COBE}, the BAO signal was measured for the first time in the CMB, detecting small angle anisotropies in the CMB angular power spectrum, confirming the cosmological predictions (\\citealt{CMBCl1}; \\citealt{CMBCl2}). Later on in 2005, the BAO signal was measured in the Sloan Digital Sky Survey \\citep[SDSS; ][]{SDSS} by \\cite{LRG} using the spatial correlation function of a spectroscopic subsample of luminous red galaxies (LRG), finding the BAO peak at $r\\approx100~h^{-1}$Mpc; and in the 2dF Galaxy Redshift Survey \\citep[2dFGRS; ][ 2003]{2dF1} by \\cite{Pk2dF} using the power spectrum. Since BAO measurements have proven to be a robust probe for extracting cosmological information, substantial work has been devoted to model and detect the BAO signal in two-point statistics and use it for cosmological constraints (e.g. \\citealt{LRGPk}; \\citealt{WPer07}; \\citealt{Sperg07}; \\citealt{BAOBOSS}; \\citealt{Ariel09}; \\citealt{Reid10}; \\citealt{Blake11}; \\citealt{Samu13}; \\citealt{wedges},b).  There are two important points related to the usual study of LSS using 3D analysis that need to be considered. First, to work in configuration-space, it is necessary to assume a fiducial cosmological model in order to transform the measured angular positions on the sky and redshifts of galaxies into comoving coordinates or distances, a process which could bias the parameter constraints if not treated carefully (see e.g. \\citealt{LRG} and \\citealt{Ariel09}). Secondly, in order to obtain a precise measurement of either the correlation function or the power spectrum, usually the whole galaxy sample is used to obtain one measurement, typically averaging over a wide redshift range assuming that the measurement at the mean redshift is representative of the entire sample, washing out information on the redshift evolution of the structures.  Even when these two issues are well understood and under control within certain conditions, a simple way to avoid them is by using two-point statistics based only on direct observables, i.e. only angular positions and/or redshifts, such as the angular correlation function $\\omega(\\theta)$ or the angular power spectrum $C_\\ell$. This is done by dividing the sample into redshift bins, or shells, in order to recover information along the line of sight, which otherwise would be lost due to projection effects. In the last few years there have been several papers modelling and analysing large galaxy catalogues using angular two-point statistics. Although most of these focus mainly on photometric-redshift galaxy surveys (\\citealt{Crocce11a}, \\citealt{Crocce11b}, \\citealt{Padm07}, \\citealt{Ross11}, \\citealt{ESanch}, \\citealt{deSimoni13}), this approach has also been applied to spectroscopic-redshift samples (\\citealt{Jacobo}, \\citealt{Jacobo2}, \\citealt{DiDio13}). Here we focus on the cosmological implications of applying this tomographic approach to a BOSS-like spectroscopic-redshift galaxy survey, computing $\\omega(\\theta)$ in redshift-shells and using this information to obtain constraints on cosmological parameters.  There are three main advantages of this tomographic approach: (i) compared to that of photometric redshifts (photo-z), the higher accuracy of spectroscopic redshifts significantly reduces the overlap between redshift-shells, allowing us to assume that there is no correlation between them due to these uncertainties, and to use thinner shells. Compared to the traditional 3D analysis, (ii) as we already mentioned, by using direct observables we do not need to assume a cosmological model in order to compute spatial separations between galaxies, their angular separations will remain unaffected independent of the cosmological model being tested; (iii) By measuring the angular scale of the BAO peak imprinted on $\\omega(\\theta)$ at many different redshifts, we are basically measuring the angular diameter distance $D_\\rmn{A}(z)$ at several redshifts instead of just one more precise measurement of the average distance $D_V(\\bar z)$ at the mean redshift of the sample, giving us more information about the rate at which $D_\\rmn{A}$ evolves, putting stronger constraints on the expansion history of the Universe.  This paper is organised as follows. In Section \\ref{sec:sample} we describe the mock catalogues that we used, how we configured them in redshift-shells, and what information we expect to extract from the technique discussed in this paper. In Section \\ref{sec:Mod} we describe the model we use for the angular correlation function measured in redshift shells and its covariance matrix, and describe the test we perform to assess the ability of these models to extract unbiased cosmological constraints. In Section \\ref{sec:BOSS} we describe a synthetic dataset constructed using these models, and present a forecast of the accuracy on cosmological constraints expected from applying this tomographic approach to the final BOSS galaxy catalogue in combination with Planck CMB measurements, comparing this with the constraints that would result from the  combination of the isotropic BAO measurements post-reconstruction on the same catalogue and Planck. We finish with our main conclusions in Section \\ref{sec:Conc}. ",
        "conclusions": "\\label{sec:Conc}  \\begin{figure} \\includegraphics[scale=0.445]{figs/w0-wa_waCDM.pdf} \\caption{The marginalised 68 and 95 per cent confidence levels in the $w_0-w_a$ plane for the extended $\\Lambda$CDM model case, with a time-dependent $w_{DE}$ parametrised as in eq. (\\ref{eq:w0wa}). The dash-dotted lines (purple contours) correspond to the constraints derived from the use of Planck+WP only. The dashed lines (green contours) are the constraints obtained by combining Planck+WP and BAO measurements post-reconstruction, while the solid lines (orange contours) are those derived from the combination of Planck+WP and $\\omega(\\theta)$ without any reconstruction. The dotted lines correspond to the fiducial values assumed to make our forecast.} \\label{fig:w0-wa} \\end{figure}  We tested the implications of applying a tomographic approach to a spectroscopic-redshift galaxy survey through measuring the two-point angular correlation function $\\omega(\\theta)$ in thin redshift shells, avoiding the need to assume a fiducial cosmological model in order to transform measured angular positions and redshifts into comoving distances, as it is the case in the traditional 3D analysis. In principle, this technique, as it is presented in this paper, can be also applied to narrow-band photometric surveys (e.g. PAU survey\\footnote{http://www.pausurvey.org/}) without any further consideration, since the accuracy in their photometric-redshifts determination is expected to be $\\sim0.0035(1+z)$, but we have not studied this case here.  We first contrasted the predictions of the model for $\\omega(\\theta)$ and its covariance matrix, described in Section \\ref{sec:Mod}, against measurements made on a set of 160 mock catalogues, and tested its ability to recover the correct value of the dark energy equation of state parameter $w_{DE}$ used to construct these catalogues. For simplicity, we did not include cross-correlations between shells in our analysis, although these should add extra information. Our modelling includes effects such as redshift-space distortions and non-linear evolution of the density fluctuations, showing that these effects are completely necessary in order to correctly reproduce the full shape of $\\omega(\\theta)$. This technique results in an unbiased way to extract cosmological information.  Next, we made a forecast of the accuracy in cosmological constraints expected from applying this technique to the final BOSS galaxy catalogue in combination with the Planck CMB results in three different flat cosmological models, and compared it with what would result from combining Planck and isotropic BAO measurements post-reconstruction on the same galaxy catalogue. To do so, we chose the best-fit of the base $\\Lambda$CDM cosmology from Planck as our true cosmology and characterised the BOSS catalogue assuming an area of $10000$ deg$^2$, a constant $n(z)=3\\times10^{-4}h^3$Mpc$^{-3}$, the redshift range $0.2<z<0.6$, and the galaxy bias model in eq. (\\ref{eq:bBOSS}). Using the model for the angular correlation function and its covariance matrix, we constructed a synthetic dataset consisting of 16 measurements of $\\omega(\\theta)$ on the same number of redshift-shells covering the whole redshift range. We also computed the CMB likelihood using distance priors for Planck, and reproduce the likelihood obtained from BAO measurements post-reconstruction on BOSS using eq. (\\ref{eq:BAO}) and assuming an error of $2\\%$ for LOWZ and $1\\%$ for CMASS.  Across this analysis, we used two different models for the galaxy bias evolution, in one case using the correct model used to construct the synthetic dataset, and in the other case using the simpler model in eq. (\\ref{eq:bLD}). The different choices showed no biasing constraining cosmological parameters, but different accuracies. The first cosmological model tested was the basic $\\Lambda$CDM, which resulted in tighter constraints for the combination of Planck and BAO measurements compared to the combination of Planck and $\\omega(\\theta)$ measurements. Although, for the other two models tested, where we extended the base model allowing $w_{DE}$ to deviate from its fiducial value of $-1$, being a constant in one case, and allowing a time-dependence in the other, we found that the more freedom we give to $w_{DE}$ the better performance our technique has, resulting in a comparable accuracy when constraining a constant $w_{DE}$ with respect to Planck+BAOs, and up to $15\\%$ higher FoM compared to the combination of Planck and BAO measurements in the case of a time-dependent dark energy equation of state, showing that this tomographic approach to analyse the galaxy clustering is able to put strong constraints in the expansion history of the Universe."
    },
    "1402/1402.1595_arXiv.txt": {
        "abstract": "NEO\u2019s come close to the Earth's orbit so that any dust ejected from them, might be seen as a meteor shower. Orbits evolve rapidly, so that a similarity of orbits at one given time is not sufficient to prove a relationship, orbital evolution over a long time interval also has to be similar. Sporadic meteoroids can not be associated with a single parent body, they can only be classified as cometary or asteroidal. However, by considering one parameter criteria, many sporadics are not classified properly therefore two parameter approach was proposed. \\vspace{-6pt} ",
        "introduction": "The association between meteoroid streams and comets is well-established, though the idea that some originate from asteroids is quite old (see Williams, 2011). The formation mechanism of a meteor stream from a comet is well-understood. Ices sublimate generating a gas outflow that carries away small dust grains.  Inter-asteroid collisions, as was seen in 2010A2 (Linear) (Jewitt {\\it et al.}, 2010) is the most obvious formation process, though there are other possibilities. In all cases, the ejection velocity is small compared to the orbital speed so that a stream is formed on an orbit that is similar to that of the parent (the physics and mathematics involved are given in Williams, 2004).  Individual meteoroid orbits can evolve in a significantly different way from the parent. The obvious mechanism causing this is a close planetary encounter (Hughes Williams \\& Fox, 1981) but collisions between stream meteoroids, the parent or interplanetary dust particles can also play a part. These processes feed the sporadic population, making it hard to associate such meteoroids with a given parent.   The major tool for determining a pairing of parent and meteoroid stream has been orbit similarity. To quantify this, Southworth \\& Hawkins (1963) and Drummond (1981) proposed orbital similarity D-criteria, both being widely used. Using such criteria, many pairings between comets and meteor streams have been established and also between minor showers and near-Earth asteroids (NEAs). Because of the vast increase in the number of known asteroids, the probability of a chance similarity of orbits is high. For this reason, Porub{\\v c}an, Korno{\\v s} \\& Williams (2004) proposed that orbits need to remain similar for 5000 years before any association can be claimed.  Three major showers, the Geminids, the Quadrantids and the Taurids all have asteroids associated with them. Whipple (1983) pointed out the orbital similarity between the Geminid stream and (3200) Phaethon, while Fox, Williams \\& Hughes (1983) showed  that Phaethon could be the parent of the Geminid meteoroid stream assuming the ejection mechanism was Cometary. Battams \\& Watson (2009) reported that Phaethon brightened by at least 2 magnitudes in 2009. The asteroids 2005 UD and 1999 YC have orbits that are similar to that of Phaethon, all suggesting that a fragmenting large comet may be the actual parent.  Jenniskens (2004) noted the similarity between the orbits of the Quadrantids % and asteroid 2003 EH1(now 196256). In the activity curve of the Quadrantids there is a broad background and a sharp narrow central peak suggesting that an old comet (probably 96P/Machholz), produced the background while a recent outburst produced the sharp peak. This could the fragmentation of C/1490 Y1 a few hundred years ago (Williams {\\it et al} 2004), producing asteroid 2003EH1.  The Taurid stream is a complex of several smaller streams and filaments. There are many associated asteroids and lists are given in Porub{\\v c}an, Korno{\\v s} \\& Williams (2006). The Taurid complex has an active comet associated with it, 2P/Encke. Asher, Clube \\& Steel (1993) suggested that the family of meteoroid streams, comet 2P/Encke and associated Apollo asteroids could all have formed by the fragmentation of a giant comet 20-30 Ky ago. \\vspace{-15pt} ",
        "conclusions": ""
    },
    "1402/1402.3629_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.2in}} There is strong evidence in favor of the idea that dark matter is self interacting, with the cross section-to-mass ratio $\\sigma / m \\sim 1~\\cm^2/\\g \\sim 1~\\text{barn}/\\gev$.  We show that viable models of dark matter with this large cross section are straightforwardly realized with non-Abelian hidden sectors.  In the simplest of such models, the hidden sector is a pure gauge theory, and the dark matter is composed of hidden glueballs with a mass around 100 MeV.  Alternatively, the hidden sector may be a supersymmetric pure gauge theory with a $\\sim 10~\\tev$ gluino thermal relic.  In this case, the dark matter is largely composed of glueballinos that strongly self interact through the exchange of light glueballs.  We present a unified framework that realizes both of these possibilities in anomaly-mediated supersymmetry breaking, where, depending on a few model parameters, the dark matter may be composed of hidden glueballinos, hidden glueballs, or a mixture of the two.  These models provide simple examples of multicomponent dark matter, have interesting implications for particle physics and cosmology, and include cases where a subdominant component of dark matter may be extremely strongly self interacting, with interesting astrophysical consequences. ",
        "introduction": "The standard model of cosmology describes a universe that is dominated by the vacuum energy $\\Lambda$ and collisionless cold dark matter (CDM). The success of the $\\Lambda$CDM model is based on its well-established record of describing the features of the large-scale structure observed in the Universe.  On smaller scales, however, the picture is much less clear.  $N$-body simulations of collisionless CDM appear to disagree with observations on small scales, motivating dark matter properties that differ significantly from the standard paradigm.  In particular, if dark matter is self interacting (able to scatter elastically with itself), simulations show that the core sizes and central densities of dwarf spheroidal galaxies, low-surface-brightness spirals, and galaxy clusters can all be brought in line with observations~\\cite{Rocha:2012jg,Peter:2012jh,Vogelsberger:2012ku,Zavala:2012us}. Modifying $\\Lambda$CDM to incorporate self-interacting dark matter (SIDM), sometimes called the $\\Lambda$SIDM model, is consistent with constraints from the Bullet Cluster, measurements of dark matter halo shapes, and subhalo survival requirements.  To make the simulations and observations consistent, the ratio of the dark matter's self-interaction cross section to its mass should be in the range $\\sigma/m \\sim 0.1 - 10~\\units{cm^2/g}$. The requirement of such strong self interactions eliminates from consideration all of the most studied dark matter candidates, including weakly interacting massive particles (WIMPs), axions, and sterile neutrinos.  At the same time, such large cross sections are on par with nuclear-scale cross sections ($1~\\units{cm^2/g} \\simeq 1.78~\\units{barn/GeV}$).  The possibility that dark matter has color and interacts through the strong interactions of the standard model (SM) is highly constrained, for example, by searches for anomalous isotopes in sea water~\\cite{Hemmick:1989ns,Verkerk:1991jf,Yamagata:1993jq}.  However, dark matter may self interact through non-Abelian forces (such as a dark analogue of QCD) in a hidden sector.  As we will show below, this is straightforwardly realized in even the simplest such hidden sectors, with SU($N$) gauge symmetry and no additional matter content. For the confinement scales $\\Lambda \\sim 100~\\mev$, the hidden gluons confine to form glueballs, and the resulting glueball dark matter has the required self interactions.  For hidden sectors that are roughly the same temperature as the visible sector, the glueball relic density is generically too large, but the desired relic density may be obtained by adjusting the relative temperatures of the hidden sector and visible sector, as we discuss below.  This hidden glueball scenario for self-interacting dark matter is remarkably simple, but it is decoupled from the visible sector, both in the technical sense and in the sense that it is not motivated by any of the well-known problems of the SM.  In addition, the correct relic density is arranged by tuning a free parameter, the ratio of hidden to visible sector temperatures, and so the scenario cannot be claimed to naturally produce the right thermal relic density, as in the case of WIMPs.  At first sight, it might appear to be difficult to enhance the model to accommodate all of these desirable features, especially since the WIMP miracle requires weak-scale annihilation cross sections, whereas the required self interactions naturally suggest strong interactions.  In fact, however, we will show that all of these features are present in a supersymmetric version of the hidden glueball scenario, in which the hidden sector is a supersymmetric pure gauge theory.  In this model, the dark matter is a $\\sim 10~\\tev$ hidden gluino, which freezes out in the early Universe when the temperature is high.  At freeze-out, the theory is weakly coupled, but, as the Universe cools and expands, the theory confines, forming hidden glueballinos and glueballs.  The glueballinos strongly interact via exchange of the hidden glueballs with the required self-interaction cross section. This scenario is straightforwardly accommodated in anomaly-mediated supersymmetry breaking (AMSB) scenarios~\\cite{Randall:1998uk,Giudice:1998xp}, which provide a connection to the problems of the SM and also allow the glueballinos to naturally inherit the correct relic density through the WIMPless miracle~\\cite{Feng:2008ya,Feng:2008mu}, a possibility discussed previously in Refs.~\\cite{Feng:2011ik,Feng:2011uf,Feng:2011in}.  For related work on strongly interacting hidden sectors and dark matter, see Refs.~\\cite{Kribs:2009fy,Alves:2009nf,Falkowski:2009yz,% Lisanti:2009am,Alves:2010dd,Kumar:2011iy,Cline:2013zca,Kang:2006yd,% Higaki:2013vuv,Heikinheimo:2013fta,Bai:2013xga}.  Of course, the supersymmetric models also contain glueballs, which, as in the nonsupersymmetric case, may be dark matter.  As we will see, in different regions of the AMSB parameter space, the dark matter may be composed of mostly glueballinos, mostly glueballs, or a mixture of the two.  For the case where the dark matter is composed of mostly glueballinos, we detail two possibilities. In the first case, the hidden sector is coupled to the visible sector only indirectly through the supersymmetry breaking mechanism. Since this coupling is extremely weak, the sectors can have different temperatures, and the glueball relic density may be very small for cold hidden sectors. An example cosmological timeline of events in this case is given in \\figref{timeline1}.  \\begin{figure}[t] \\centering \\includegraphics[scale=0.7]{running} \\caption{Example timeline of events in the supersymmetric pure SU($N$) theory without connectors, in terms of the hidden- and visible-sector temperatures $T^h$ and $T$.  The hidden-sector coupling $\\alpha_h$ is shown as a function of these temperatures. It is weak at gluino freeze-out but grows as the temperature drops, leading to confinement and the formation of glueballino and glueball dark matter at a temperature $\\sim \\Lambda$. The scenario is described in detail in \\secref{noconnectors}, and the chosen parameters are represented by the yellow dot in \\figref{AMSB-pure-relic}.  } \\label{fig:timeline1} \\end{figure}  In the second case, the hidden sector is coupled to the visible sector through connector fields.  The visible and hidden sectors, therefore, have the same temperature at early times, leading, {\\em a priori}, to too-large glueball relic densities.  Decays of glueballs are in conflict with big bang nucleosynthesis (BBN) and other astrophysical and particle constraints.  Instead, we rely on a novel nonthermal process in the early Universe to deplete the gluon density, thereby suppressing the glueball density after confinement.  In this case, the gluons annihilate into singlet right-handed neutrinos with $\\sim 1~\\units{GeV}$ mass, and we reduce the hidden gluon density by forcing the right-handed neutrinos to decay into SM particles more quickly than they can annihilate back into hidden gluons.  A representative timeline for this case is shown in \\figref{timeline2}.  \\begin{figure}[t] \\centering \\includegraphics[scale=0.7]{running-withConn} \\caption{As in \\figref{timeline1}, but for supersymmetric pure SU($N$) theory {\\em with} connectors.  Since the hidden and visible sectors communicate efficiently in the early Universe, they share a temperature, $T$.  The gluon population is depleted through annihilations to and the subsequent decays of the $\\nu_R$ in the visible sector, and the resulting scenario has pure glueballino dark matter.  The scenario is described in detail in \\secref{connectors}, and the chosen parameters are represented by the yellow dot in \\figref{AMSB-pureConn-relic}. } \\label{fig:timeline2} \\end{figure}  This work is organized as follows. In \\secref{evidence}, we review the astrophysical evidence for self-interacting dark matter.  In \\secref{glueball}, we begin with the simplest possible case: nonsupersymmetric pure gauge hidden sectors and glueball dark matter. We discuss glueball self interactions and relic densities and determine the preferred parameters for this simple model.  We then move to supersymmetric models with pure gauge hidden sectors and glueballino dark matter.  In \\secref{glueballinoscattering}, we review the calculation of the glueballino self-interaction cross section, and in \\secref{glueballinorelic}, we discuss the glueballino relic density and the realization of the WIMPless miracle in the AMSB framework. Finally, with this groundwork, we present full AMSB models of glueballino/glueball dark matter without and with connectors in \\secsref{noconnectors}{connectors}, respectively. We conclude in \\secref{conclusions}.  Last, we make a note on naming conventions. In the rest of this work, we follow the literature: glueballinos denote gluino-gluon bound states, while gluinoballs denote gluino-gluino bound states.  In addition, unless otherwise stated, gluon, gluino, glueball, and glueballino refer to hidden sector particles and are denoted by $g$, $\\tilde{g}$, $\\gb$, and $\\gbino$, respectively. ",
        "conclusions": "\\label{sec:conclusions}  We have explored the possibility that dark matter may be a composite particle, made up of bound states of a dark analogue of QCD in the hidden sector.  Such constructions lead to rich and varied phenomena that are distinct from the WIMP scenario more typically considered. It also naturally leads to large self interactions of the dark matter, which can explain several observational puzzles in the small-scale structure of the Universe.  The simplest scenarios contain only dark gluons, which confine into glueballs with cosmologically interesting scattering cross sections for confinement scales around 100 MeV.  Arranging the correct relic density requires one to disconnect the temperatures in the hidden and visible sectors such that their ratio at confinement is $\\sim 10^{-3}$.  A richer theory arises when one considers supersymmetric versions, for which the dark gluino mass provides a separate mass scale and (in AMSB) can provide the correct relic density of glueballinos via the WIMPless miracle.  The phenomenology depends crucially on how connected the hidden sector is to the visible matter.  If there are no light connecting particles, one can dial the balance of dark matter from glueballs to glueballinos by adjusting the relative temperatures of the hidden and visible sectors.  These mixed scenarios are strongly interacting analogues of atomic dark matter~\\cite{Kaplan:2009de,CyrRacine:2012fz,Cline:2013pca,Kahlhoefer:2013dca,Fan:2013yva} and inspire further simulation of the galactic dynamics in cases where there are two components of dark matter with naturally very different mass scales and different self-interaction rates.  Such simulations would be very helpful to better understand the observational limits on these theories.  For clusters, another important issue is the fact that the dark matter may have enough energy to scatter inelastically, bringing the details of the dark composite sector to the forefront of the physics; further work is needed to better understand the implications. We have also pointed out that our models have rich implications for the early growth of supermassive black holes. The mechanism by which $\\sim 10^9 \\, {\\rm M}_\\odot$ quasars are assembled as early as redshifts of 6--7 is a mystery, and self-interacting dark matter could have a major role to play in this story.  If the hidden and visible sectors are closely connected such that the temperatures remain comparable even at late times, the hidden glueballs will generically overclose the Universe.  We considered a depletion mechanism into right-handed neutrinos and found that it can efficiently remove hidden gluons before confinement.  Self-interaction strengths required to explain the astrophysical puzzles on small scales are obtained for glueballino masses $\\gtrsim 1$~TeV and confinement scales $\\sim$~MeV.  The possibility of strong self interactions in the dark sector is well motivated by observations of lower-than-expected dark matter densities in the centers of galaxies. A strongly interacting hidden sector naturally realizes this possibility. Even in the simple models explored in this paper, we have discovered new features that must be incorporated into numerical simulations to correctly predict the spatial distribution of dark matter in the central parts of structures from dwarf galaxies to clusters of galaxies."
    },
    "1402/1402.1901_arXiv.txt": {
        "abstract": "Helioseismology provides important constraints for the solar dynamo problem. However, the basic properties and even the depth of the dynamo process, which operates also in other stars, are unknown. Most of the dynamo models suggest that the toroidal magnetic field that emerges on the surface and forms sunspots is generated near the bottom of the convection zone, in the tachocline. However, there is a number of theoretical and observational problems with justifying the deep-seated dynamo models. This leads to the idea that the subsurface angular velocity shear may play an important role in the solar dynamo. Using helioseismology measurements of the internal rotation and meridional circulation, we investigate a mean-field MHD model of dynamo distributed in the bulk of the convection zone but shaped in a near-surface layer. We show that if the boundary conditions at the top of the dynamo region allow the large-scale toroidal magnetic fields to penetrate into the surface, then the dynamo wave propagates along the isosurface of angular velocity in the subsurface shear layer, forming the butterfly diagram in agreement with the Parker-Yoshimura rule and solar-cycle observations. Unlike the flux-transport dynamo models, this model does not depend on the transport of magnetic field by meridional circulation at the bottom of the convection zone, and works well when the meridional circulation forms two cells in radius, as recently indicated by deep-focus time-distance helioseismology analysis of the SDO/HMI and SOHO/MDI data. We compare the new dynamo model with various characteristics if the solar magnetic cycles, including the cycle asymmetry (Waldmeier's relations) and magnetic `butterfly' diagrams. ",
        "introduction": "The dynamo which operates in the solar convection zone and defines the properties of sunspot cycles remains enigmatic despite substantial observational and modeling efforts. In particular, the solar dynamo models must explain the magnetic \"butterfly\" diagram (Fig.~\\ref{fig1}) and asymmetry of the sunspot cycles (Fig.~\\ref{fig2}a), so-called Waldmeier's effect \\citep{Waldmeier1935}. The magnetic butterfly diagram is obtained by stacking azimuthally averaged synoptic magnetic field maps, provided by the National Solar Observatory since 1976.  \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.65]{fig1a.eps} \\caption{Magnetic `butterfly' diagram shows the evolution of the azymuthally averaged line-of-sight magnetic field a function of latitude and time. The range of magnetic field strength is from -10~G to +10~G. The synoptic magnetogram data are provided by the National Solar Observatory \\label{fig1}} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.5]{fig2.eps} \\caption{a) The sunspot number as a function of time illustrates the asymmetry of the solar cycles (the Waldmeier's effect); b) comparison of the sunspot number curves for four cycles \\citep{Bracewell1988}; c) a model of the Waldmeier relations \\citep{Bracewell1988}; d) dependence of the cycle amplitude from the growth time \\citep{Kitiashvili2009}. \\label{fig2}} \\end{center} \\end{figure}  The most prominent features of the butterfly diagram are migration of the sunspot formation zones towards the equator, migration of the following polarity flux towards the poles, polarity reversals of polar magnetic fields during sunspot maxima, and also polarity reversals of the toroidal magnetic field during the solar minima, which are observed as cyclic changes of polarity of leading and following sunspots in bipolar active regions (the Hale's law). The Waldweier's effect (non-linear asymmetry of sunspot cycles) has the following three main characteristics: 1) the sunspot number growth time is shorter than the decay time (Fig.~\\ref{fig2}b); 2) the growth time of strong cycles is shorter than the growth time of weak cycles (Fig.~\\ref{fig2}c-d);  3) the strong cycles are shorter than the weak cycles. The first property is often used for predicting the sunspot maximum from the growth rate at the beginning of a cycle, using the Waldmeier's \"standard curves\" (Fig.~\\ref{fig2}c). Thus, the dynamo theories have to explain both the Hale's law and the Waldmeier's effect. ",
        "conclusions": "The helioseismology discoveries of the rotational subsurface shear layer and the double-cell meridional circulation require to re-examine dynamo models of the solar activity cycles. They lead a new paradigm of the solar dynamo distributed in the convection zone and equator-ward migrating dynamo waves in the sub-surface shear layer. The double-cell meridional circulation affects the dynamics of the large-scale magnetic fields, and puts additional constraints. In particular, if the observed equator-ward migration of the sunspot formation zone (`butterfly' diagram) is linked to migration of toroidal magnetic field in the convection zone then the magnetic field of sunspots may emerge from the depth of about 120~Mm. This is generally consistent with the observed rotation rate of surface magnetic fields of active regions and the internal differential rotation determined by helioseismology. Further helioseismology investigation will shed more light on the solar dynamo mechanism."
    },
    "1402/1402.0888_arXiv.txt": {
        "abstract": "We investigate the properties and clustering of halos, galaxies and blackholes to $z=0$ in the high resolution hydrodynamical simulation MassiveBlack-II (MBII). MBII evolves a $\\Lambda$CDM cosmology in  a cubical comoving volume of $V_{box} = (100 {\\rm Mpc/h})^3$ and is able to resolve halos of mass $\\mhalo = 10^9\\msun$/h. It is the highest resolution simulation of this size which includes a self-consistent model for star formation, black hole accretion and associated feedback. We provide a simulation browser web application which enables interactive search and tagging of halos, subhalos and their properties and publicly release our galaxy catalogs to the scientific community. Our analysis of the halo mass function in MBII reveals that baryons have strong effects, with changes in the halo abundance of  20-35\\% below the knee of the mass function ($\\mhalo \\lsim 10^{13.2}\\msun$/h at $z=0$) when compared to fits based on dark matter only simulations. We provide a fitting function for the halo mass function valid for the full range of halo masses in MBII out to redshift $z=11$ and discuss how the onset of non-universal behavior in the mass function limits the accuracy of our fit. We examine  the halo occupation distribution of satellite galaxies and present results valid over 5 orders of magnitude in host halo mass. We study the clustering of galaxies, and in particular the evolution and scale dependence of stochasticity and bias. Comparison with observational data for these quantities for samples with different stellar mass thresholds yields reasonable agreement. Using population synthesis, we find that the shape of the cosmic spectral energy distribution predicted by MBII is consistent with observations, but lower in amplitude. The Galaxy Stellar Mass Function (GSMF) function is broadly consistent with observations at $z \\ge 2$. At $z < 2$, observations probe deeper into the faint end and the population of passive low mass (for $M_{*}<10^{9} \\msun$) galaxies in the simulation makes the GSMF too steep. At the high mass end ($M_{*}>10^{11} \\msun$) galaxies hosting bright AGN make significant contributions to the GSMF. The quasar bolometric luminosity function is also largely consistent with observations. We note however that more efficient AGN feedback (beyond simple thermal coupling used here) is likely necessary for the largest, rarest objects/clusters at low redshifts. ",
        "introduction": "The Cold Dark Matter model with a cosmological constant ($\\Lambda$CDM) is well established enough (see e.g.,  \\citealt{2013arXiv1303.5076P},\\citealt{2013ApJS..208...19H}, \\citealt{2014arXiv1401.1389L},\\citealt{2013arXiv1311.0840C}) that individual  large-scale simulation efforts can be carried out that focus on just this one cosmology. We have also reached the point at which supercomputers enable numerical modeling of cosmological volumes with enough resolution to study the properties of individual galaxies. In this paper we report on a {\\small P-GADGET} hydrodynamic simulation of $100 \\hmpc$ cubic volume, the MassiveBlackII simulation. It has $\\sim 10^{6}\\msun$ mass resolution, cooling, star formation, black holes and feedback, and represents the evolution of a $\\Lambda$CDM universe to redshift $z=0$.   Numerical simulations  (see reviews by \\citealt{2008SSRv..134..229D}, \\citealt{2012AN....333..515S}, \\citealt{1998ARA&A..36..599B}) are the tool of choice to address many questions in cosmology, as galaxy formation is a complex non-linear problem. Two criteria which must be satisfied for accurate results are:  (1) A large enough simulation volume that Fourier density modes on the largest scales are evolving independently. The volume simulated must be a representative region of the Universe, otherwise inferences drawn from it regarding such questions as clustering and the mass function of objects will be incorrect (\\citealt{2005MNRAS.358.1076B}, \\citealt{2006MNRAS.370..993B}, \\citealt{2013arXiv1312.5256O}). When $\\Lambda$CDM models are evolved to redshift $z=0$, then a volume of at least $\\sim 100 \\hmpc$ on a side becomes necessary to predict the overall star formation rate, for example (\\citealt{2003MNRAS.339..312S}).    (2) High enough mass and spatial resolution that the properties of the objects of interest have converged. This requires many resolution elements (particles or grid cells). If we focus on particle based simulations relevant to the current work, at the very least for identification of objects there must enough particles to overcome shot noise. If we require detailed properties of the objects such as galaxy spectra, or angular momenta then this can require many more ( e.g., \\citealt{2007MNRAS.374.1479G}).   These criteria of large simulation volume and high resolution are more straightforward to address in the context of more restricted physical modeling. As a result,  dark matter and gravity-only simulations have long been used to make cosmological predictions that cover both large and small scales in the same volume (e.g., \\citealt{2005Natur.435..629S}, \\citealt{2009MNRAS.398.1150B}, \\citealt{2011ApJ...740..102K}). Semi-analytic modeling has been used to process dark matter simulations, resulting in many studies of the galaxies and their properties in the $\\Lambda$CDM model (see e.g., \\citealt{2006RPPh...69.3101B}, \\citealt{2012MNRAS.419.3200H} and references therein).  Baryonic physics including hydrodynamics obviously plays an important role in the formation of luminous objects and structure. This has lead to the inclusion of the relevant equations in simulation codes in many forms, Smoothed Particle Hydrodynamics (SPH, \\citealt{1992ARAA..30..543M}), Eulerian grid solvers (e.g., \\citealt{1992ApJS...78..341C}, \\citealt{1997ASPC..123..363B}) and hybrid Lagrangian/Eulerian schemes (e.g., \\citealt{2010MNRAS.401..791S}). Although previous work has not simultaneously reached the combination of large volume and high resolution that we present here, research has progressed using many methods, including making use of zoom simulations of smaller volumes inside a representative one (\\citealt{1993ApJ...412..455K}, \\citealt{2012MNRAS.423.1726S}), simulations that stop at high redshifts before large-scale modes become non-linear (e.g., \\citealt{2012ApJ...745L..29D}), or by tackling problems which require lower mass resolution (e.g., \\citealt{2012ApJ...758...74B}).  The advent of large-scale computing facilities with 100,000 compute cores or more (such as the Cray XT5, ``Kraken'' on which the current simulation was run) and the development of highly efficient  distributed memory simulation codes  (such as {\\small GADGET2}, \\citealt{2005MNRAS.364.1105S}) means that simulations which satisfy both criteria (1) and (2) are now possible. We have run one such simulation as part of the NSF Petascale Applications in Cosmology program, using the code {\\small P-GADGET} (see e.g., \\citealt{2012ApJ...745L..29D}).  Our aim was to simulate and analyze a large, representative volume of the $\\Lambda$CDM model with the most important physical processes previously included in zoom runs or simulations with smaller boxes. These are hydrodynamics (using SPH), cooling, a subgrid multiphase model for star formation (\\citealt{2003MNRAS.339..289S}) and subgrid black hole modeling (\\citealt{2005MNRAS.361..776S},\\citealt{2005Natur.433..604D}), both with feedback. Our use of the physical modeling and algorithms used in previous work such as \\cite{2008ApJ...676...33D}, \\cite{2009MNRAS.400...43C}, \\cite{2010MNRAS.402.1927D} and \\cite{2012ApJ...745L..29D} enables continuity and therefore comparison with this previous work. Our aim is to see what this ``fiducial'' model ($\\Lambda$CDM + {\\small GADGET} SPH + the particular subgrid  algorithms employed) predicts about the properties of galaxies, their halos and their clustering at redshifts extending down to the present day. We have not adjusted methods, algorithms and parameters used in previous work (e.g., \\citealt{2012MNRAS.424.1892D}) to try to tune to observational results. Our goal is to see how this model performs, now that there is a large volume at high resolution. We naturally expect both regions of agreement and disagreement with observations and we aim that our work  will offer guidance to future work to address the problems.  In this paper, we make the first use of the MassiveBlackII simulation evolved to $z=0$, and  use it to explore some topics in structure formation. Here we choose to focus on topics relevant to galaxy and AGN formation and large-scale structure, including mass functions, galaxy and halo properties and clustering. Our emphasis is on lower redshifts; the simulation at redshifts $z>5$ has been explored by \\cite{2013MNRAS.430.2885W}, \\cite{2013MNRAS.429.2098W}, and \\cite{2013MNRAS.435.2885W}. Topics that we leave to future work include the intergalactic medium, absorption lines, galaxy clusters and X-ray emission.  Our plan for the paper is as follows. In Section~\\ref{sec_methods} we briefly describe the numerical methods and algorithms used to run the simulation, select galaxies and carry out stellar population synthesis. In Section~\\ref{sec_vis}, we describe visualization of the simulation. In Sections \\ref{sec_massfn} and \\ref{sec_hod} respectively we present the mass function  and halo occupation distribution and in Section \\ref{sec_gal_clustering} we examine the clustering of dark matter and galaxies. The properties of galaxies and supermassive black holes are examined in Sections~\\ref{sec_galaxy}  and \\ref{sec_bh} and we derive some conclusions from our work in Section~\\ref{sec_conc}. ",
        "conclusions": "\\label{sec_conc}  In this paper we have examined a variety of standard predictions using the recently completed MBII simulation. In particular, the results presented here range from the clustering of halos and galaxies, HOD, the mass function of halos to basic properties of galaxies and AGNs.  Our main conclusions are:  \\begin{itemize}  \\item We find that halos with masses $M_{halo} \\geq 10^{13}\\msun/h$ have a baryon fraction close to 80\\%-90\\% of the cosmic mean. The baryon fraction decreases steadily with decreasing halo mass. Our results are broadly consistent with \\cite{2007MNRAS.377...41C} for cluster sized halos but differ significantly for lower mass halos. The discrepancy can be attributed to feedback both from AGN and supernovae which were not been modeled in \\citet{2007MNRAS.377...41C}.  \\item We find that the FOF halo mass function (where halo mass includes both dark matter and baryonic components) in MBII can be fit with a universal form (valid for all redshifts) at the $\\sim 13\\%$ level.  \\item One of the most striking results predicted by MBII is the behavior of the halo mass function which shows a strong suppression in the abundance of halos below the knee of the mass function ($\\ln(1/\\sigma) \\lsim -0.2$ which corresponds to $\\mhalo \\sim 10^{13.2}\\msun$/h at $z=0$ and $\\mhalo \\sim 10^{9.5}\\msun$/h at $z=3$) when compared to dark matter only predictions of the halo mass function. This is due to the significant impact of baryonic processes which tend to suppress the mass of the halo by up to 30\\% \\cite{2013MNRAS.431.1366S}.  \\item Fits to the mass function from dark matter only simulations overpredict the  mass function at the 20-35\\% level below the knee of the mass function at $\\ln(1/\\sigma) \\lsim 0.2$.  \\item We have quantified the scale dependence of bias and stochasticity in the simulation. We find that scale of the deviation from a linear fit for both bias and stochasticity reaches a minimum at $z\\sim 3-4$. The galaxy and mass density field can in this sense be said to trace each other most closely at this redshift.  \\item The MBII overall matches most observed measurements of galaxy clustering with stellar mass. There are however some discrepancies in the bias at the lowest masses and redshifts probed. The simulation is consistent with observed measurements of stochasticity from combinations of galaxy clustering and weak lensing.  \\item We find that the HOD is well described by a power law behavior (see equation~\\ref{OccFit}). We find a modest evolution for the power law slope however the normalization mass for the distribution exhibits an exponentially decaying behavior with redshift.  \\item The location of the peak and the width for the radial distribution of satellite galaxies decrease with decreasing redshift irrespective of the host mass of the halo, showing that satellite galaxies cluster more strongly around the central galaxy with time.  \\item We find that the shape of the CSED in MBII is consistent with observed data. The amplitude is however much lower and can be attributed to incompleteness at low galaxy masses in MBII.  \\item The GSMF predicted by MBII is consistent with observations out to $z=2$. At lower redshifts however the  MBII GSMF is much steeper at lower masses whereas MBII over produces stars at larger masses at $z=0$.  \\item We find black hole mass and AGN luminosity to be broadly correlated to host $M_{\\rm halo}$ but with very large scatter, indicating a wide range of black hole properties for a given halo mass.  However, the local $M_{\\rm BH} - M_{*}$ relation is tighter and consistent the observed one. Our results also imply very moderate evolution for the $M_{\\rm BH} - M_{*}$ from $z\\sim 2$.  The global star formation rate density and black hole accretion rate density are also similar, but with a peak for the latter shifted to later times.  \\item The bolometric luminosity function of the simulated AGN population is broadly consistent with observational constraints \\citep[][see also]{2013ApJ...773...14R}. We note however that simple thermal coupling for AGN feedback appears to quench the bright end of the AGN LF too fast around $z \\sim 2$ whilst not being sufficient to fully quench the brightest AGN at $z=0$. The best agreement with observations is at $z > 2$ and $z \\sim 1$. Our results show a weak dependence of AGN clustering with luminosity. (Note that we defer to future work we to carry out detailed analysis of AGN population). \\end{itemize}  The MBII simulation is the largest simulation of its kind run to date with sufficient resolution to resolve $10^9\\msun$/h halos.  We have found that the properties, such as HOD and clustering are consistent with previous work and observations. However we find that feedback from AGNs is still not sufficient at lower redshifts to properly account for the properties of large galaxies and AGNs since it is unable to quench star formation in massive galaxies and the abundance of luminous quasars.  However it describes the current state of standard SPH simulations of galaxy formation and should be used as a testbed for improving models of galaxy formation.  The parameters used in MBII were based on much smaller SPH simulations which missed these large halos in their volumes and thus could not find discrepancies with observations at the large mass end.  At smaller masses the GSMF in MBII is very steep compared to observational constraints. A variable wind model based on the galaxy velocity dispersion \\citep{2006MNRAS.373.1265O,2011MNRAS.415...11D, 2013MNRAS.436.3031V, 2014MNRAS.tmp...38T} has been shown to flatten the GSMF at smaller masses to better reproduce observational constraints at $z=0$. However at higher redshifts these models still predict a steep slope. On the other hand we find that AGN feedback is still insufficient to quench star formation in the most massive galaxies. The effect of AGNs is important in regulating the growth of massive galaxies; e.g. the results of \\citet{2011MNRAS.415...11D} (which do not include AGN feedback)  and those of \\citet{2013MNRAS.436.3031V, 2014MNRAS.tmp...38T} (which include AGN feedback) clearly illustrate how AGNs affect the stellar content of massive galaxies. However these simulations are still small in volume and miss out on massive halos. It would be interesting to see their predictions when run on boxes similar in size to MBII. In subsequent work we will look at improving on the feedback models employed in MBII.  One of the important effects we have quantified here is the effect of baryonic processes on the halo mass function. Although we find that the universality of the mass function in MBII holds at the 13\\% level consistent with previous work, we find considerable differences in the abundance of halos below the knee of the mass function when compared to dark matter simulations. One of the natural questions which arise are: how do different models for star formation, black hole growth and feedback affect the mass function? How strong is the redshift dependence? What can be said of halo abundance matching techniques and their predictions when done with mass functions such as MBII? These questions can only be addressed with simulations with different galaxy formation models and corresponding dark matter simulations  which we will address in future work."
    },
    "1402/1402.5463_arXiv.txt": {
        "abstract": "Using 72 Short Gamma Ray Bursts (SGRBs) with well determined spectral data observed by BATSE, we determine their redshift and the luminosity by applying $E_p$--$L_p$ correlation for SGRBs found by \\cite{tsutsui13}. For 53 SGRBs with the observed flux brighter than $4 \\times 10^{-6}~{\\rm erg~cm^{-2}s^{-1}}$, the cumulative redshift distribution up to $z=1$ agrees well with that of 22  {\\it Swift}~SGRBs. This suggests that the redshift determination by the $E_p$--$L_p$ correlation for SGRBs works  well. The minimum event rate at $z=0$ is estimated as $\\rho_{SGRB}(0) = 6.3_{-3.9}^{+3.1} \\times 10^{-10}~{\\rm events~Mpc^{-3}yr^{-1}}$ so that the minimum beaming angle is $0.6^\\circ-7.8^\\circ$ assuming the merging rate of $10^{-7}-4\\times 10^{-6}~{\\rm events~Mpc^{-3}yr^{-1}}$ suggested from the binary pulsar data. Interestingly, this angle is consistent with that for SGRB130603B of $\\sim 4^\\circ-8^\\circ$\\citep{fong13b}. On the other hand, if we assume the beaming angle of $\\sim 6^\\circ$ suggested from four SGRBs with the observed value  of beaming angle, the minimum event rate including off-axis SGRBs is estimated as $\\rho_{SGRB,all}^{min}(0)=1.15_{-0.71}^{+0.57}\\times 10^{-7}~{\\rm events~Mpc^{-3}yr^{-1}}$. If SGRBs are induced by coalescence of binary neutron stars (NSs) and/or black holes (BHs), this event rate leads to the minimum gravitational-wave detection rate of $\\rm 3.9_{-2.4}^{+1.9} (152_{-94}^{+75})~events~y^{-1}$ for NS-NS (NS-BH) binary, respectively, by a worldwide network with KAGRA, advanced-LIGO, advanced-Virgo, and GEO. ",
        "introduction": "Although the number of observed Short Gamma-Ray Bursts (SGRBs) is increasing, their central engine is still a big mystery. A major candidate is coalescence of compact objects such as neutron stars and stellar-mass black holes. One of the keys to confirm this idea is the formation rate of SGRBs as a function of redshift. In fact, if SGRBs are truly induced by coalescence of compact objects, the SGRB formation rate will track the star formation rate with some delay time. Further, if this is the case, SGRBs are expected to be accompanied by substantial gravitational-wave emission. Thus, the local SGRB formation rate is directly related to the expected number of gravitational-wave events for the next-generation gravitational-wave detectors such as KAGRA \\footnote{http://gwcenter.icrr.u-tokyo.ac.jp/en/}, advanced-LIGO \\footnote{http://www.ligo.caltech.edu/}, advanced-VIRGO \\footnote{http://www.ego-gw.it/index.aspx/} and GEO \\footnote{http://www.geo600.org/}.  However, the number of SGRBs with known redshift is very small ($\\sim 20$) so that the formation rate is not easy to estimate. Previous studies have estimated the formation rate assuming the functional form of the formation rate and the luminosity function and fitting the data to derive model parameters \\citep{guetta05,guetta06,nakar06,dietz11,metzger12,petrillo13}. In this approach, the result depends on the functional form of the model and has large statistical errors due to the small number of SGRBs used to fit the model.  On the other hand, as to Long Gamma-Ray Bursts (LGRBs), the formation rate has been estimated much more precisely and robustly. This is because the correlation between the spectral peak energy and luminosity was found and used to estimate the redshift of LGRBs without known redshift. First, \\cite{yonetoku04} analyzed data of 16 LGRBs observed by {\\it CGRO} BATSE and {\\it BeppoSAX} with known redshifts, and found the $E_p-L_p$ correlation between the $E_p$ and the peak luminosity $L_p$ integrated for 1~s time interval at the peak as, \\begin{equation}\\label{eq:lgrbeplp} L_p=\\rm 2.34\\times10^{51}erg~s^{-1}(\\frac{E_p}{100keV})^2. \\end{equation} The linear correlation coefficient of the $\\log E_p-\\log L_p$ correlation is 0.958 and the chance probability is $5.31 \\times 10^{-9}$. Then \\cite{ghirlanda05a, ghirlanda05b, krimm09, yonetoku10} checked the properties of the correlation and confirmed its reliability. Using this correlation, \\cite{yonetoku04} estimated the redshift for 689 bright BATSE LGRBs without known redshift and derived the luminosity function and the formation rate.  As to SGRBs, however, due to the small number of the events with known redshifts and the good spectra to determine $E_p$, it has been difficult to do a similar analysis. Recently, \\cite{tsutsui13} succeeded in determining $E_p-L_p$ correlation for SGRBs. They used 8 secure SGRBs out of 13 candidates and obtained \\begin{equation}\\label{eq:sgrbeplp} L_p = \\rm 7.5 \\times 10^{50}erg~s^{-1}(\\frac{E_p}{100keV})^{1.59}, \\end{equation} where $E_p$ is from the time integrated spectrum again while $L_p$ was taken as the luminosity integrated for 64~ms time intervals at the peak considering the shorter duration of SGRB. The linear correlation coefficient of the $E_p-L_p$ correlation is 0.98 and the chance probability is $1.5 \\times 10^{-5}$. Although this is not so tight as that for LGRBs due to the fact that the number of SGRBs is half of that of LGRBs, it is accurate enough to use as a redshift indicator for many SGRB events without known redshifts.  In this article, we determine the redshifts of SGRBs observed by BATSE using the $E_p$--$L_p$ correlation mentioned above. Then, we obtain a non-parametric estimate of the luminosity function and SGRB formation rate versus redshift based on much more samples compared with previous studies. This article is organized as follows. In \\S 2 , we describe the observations and data analyses. After that, we show the redshifts estimated by $E_p$--$L_p$ correlation for SGRBs, and obtain the cumulative redshift distribution and compare it with the observed one. We also show the cumulative luminosity function and the SGRB formation rate as a function of redshift with non-parametric method, i.e. without any assumptions on both distributions. \\S 3 is devoted to discussions and implications of the results. ",
        "conclusions": "\\label{discussion} \\label{sec:discussion}  Long GRBs(LGRBs) are believed to be caused by relativistic jets since the break of the afterglow light curves are seen for many LGRBs. The typical example is GRB990510 which shows achromatic break of the afterglow light curve \\citep{Harrison1999}. The physical reason for the achromatic break of the light curve comes from the jet dynamics after the relativistic $\\Gamma$ factor of the jet becomes smaller than the inverse of the jet opening angle \\citep{Rhoads1999}. If SGRBs are also caused by relativistic jets, we can expect a jet break similar to that of LGRBs. However, so far, only GRB~130603B has multi-wavelength data from radio to X-ray to confirm the achromatic jet break of the afterglow. \\citet{fong13b} determined the jet opening angle of GRB~130603B as $4^\\circ - 8^\\circ$, where the ambiguity comes from the uncertainty in the kinetic energy of the jet and the ambient gas density. Very interestingly, the opening angle of GRB~130603B is compatible with the minimum jet opening angle derived in the previous section.  For GRB~051221, the steepening of the afterglow light curve was observed in X-ray band at $t \\sim 5$ days. If we identify this steepening as the jet break, the jet opening angle is determined as $5.7^\\circ -7.3^\\circ$ \\citep{soderberg06}. For GRB111020A, \\citet{fong12} argued a significant break at $t \\sim 2$ days in X-ray band and obtained the jet opening angle of $3^\\circ - 8^\\circ$. For GRB090426, \\citet{nicuesa11} obtained the jet opening angle of $4.4^\\circ$ from a break at $t = 0.4$ day. The other jet opening angle information from SGRBs is lower limits (see section 8.4 of recent review by \\citet{berger13}) so that we use only these four estimations of the jet opening angles. A simple average of these four angles is $\\sim 6^\\circ$. Taking this value, then, the event rate of SGRB including the off-axis ones becomes $\\rho_{SGRB, all}^{min}(0)= 1.15_{-0.71}^{+0.57}\\times 10^{-7}~{\\rm events~Mpc^{-3}yr^{-1}}$.   Now let us assume that the central engine of SGRB is a coalescence of a neutron star-neutron star (NS-NS) binary. Then, we can obtain an estimate of the event rate of gravitational-wave detection. The detectable range of KAGRA, adv-LIGO, adv-VIRGO and GEO network is $\\rm \\sim 200~Mpc$ so that the minimum gravitational-wave detection rate is obtained as $\\rm 3.8_{-2.2}^{+1.8} ~events~y^{-1}$. This is an estimate independent of the one based on pulsar observations \\citep{Kalogera04a, Kalogera04b}. On the other hand, if the central engine of SGRBs is a coalescence of a black hole-neutron star (BH-NS) binary with masses, say $M_{\\rm BH}= 10 M_{\\odot}$ and $M_{\\rm NS}=1.4 M_{\\odot}$, the detectable range will become $\\sim 3.4$ times larger because it is proportional to $M_{\\rm chirp}^{5/6}$, where the chirp mass $M_{\\rm chirp}$ of a binary is defined by $M_{\\rm chirp} = (M_1M_2)^{3/5}/(M_1+M_2)^{1/5}$ \\citep{Seto2001} where $M_1$ and $M_2$ are the masses of each compact object. In this case, the detection rate will be $\\rm 146_{-83}^{+71}~events~y^{-1}$ .  In the above estimation, we used the BATSE SGRBs with a flux larger than $4 \\times 10^{-6}~{\\rm erg~cm^{-2}s^{-1}}$ which is shown by the solid line in figure~\\ref{fig1}.  As is seen in figure~\\ref{fig1}, there are 17 events below the flux limit and the event with the lowest fluxes located a factor of $\\sim 4$ below the solid line. If we lower the flux limit by a factor of 4, noting that the cumulative luminosity function is proportional to $L^{-1}$ as seen in figure~\\ref{fig4}, we can expect that the event rate can be a factor $\\sim 4$ larger. Then the gravitational-wave detection rate becomes $\\rm 15.2_{-8.8}^{+7.2} (584_{-332}^{+284})~events~y^{-1}$  for NS-NS (NS-BH) binary, respectively. While \\cite{abadie10}  suggested that the likely binary neutron-star detection rate for the Advanced LIGO-Virgo network will be 40 events per year, with a range between 0.4 and 400 per year in their review paper.  \\citet{coward12} analyzed 14 {\\it Swift} SGRBs taking into account the ratio of the rate of BATSE to {\\it Swift} SGRBs, the k-correction, the maximum distance observed by a certain event, the jet opening angle and the probability of the event being SGRB. They obtained the event rate of $8_{-3}^{+5}\\times 10^{-9}$--$1.1_{-0.47}^{+0.7}\\times 10^{-6} ~{\\rm events~Mpc^{-3}yr^{-1}}$ for the case of isotropic emission and beamed emission, respectively. In our analysis the minimum event rate is $\\rho_{SGRB, all}^{min}(0)=1.15_{-0.71}^{+0.57}\\times 10^{-7}~{\\rm events~Mpc^{-3}yr^{-1}}$ after the geometrical correction of beaming angle, while the realistic one is $4.60_{-2.84}^{+2.28}\\times 10^{-7}~{\\rm events~Mpc^{-3}yr^{-1}}$. Although both analysis used completely different methods, they are consistent with each other. For NS-NS binary, the detection range of adv-LIGO will be $\\rm \\sim 100~Mpc$ in 2016-2017 \\citep{abadie10,aasi13} so that there is a good chance of the first gravitational wave detection around 2017."
    },
    "1402/1402.5978_arXiv.txt": {
        "abstract": "We present the use of continuous-time autoregressive moving average (CARMA) models as a method for estimating the variability features of a light curve, and in particular its power spectral density (PSD). CARMA models fully account for irregular sampling and measurement errors, making them valuable for quantifying variability, forecasting and interpolating light curves, and for variability-based classification. We show that the PSD of a CARMA model can be expressed as a sum of Lorentzian functions, which makes them extremely flexible and able to model a broad range of PSDs. We present the likelihood function for light curves sampled from CARMA processes, placing them on a statistically rigorous foundation, and we present a Bayesian method to infer the probability distribution of the PSD given the measured lightcurve. Because calculation of the likelihood function scales linearly with the number of data points, CARMA modeling scales to current and future massive time-domain data sets. We conclude by applying our CARMA modeling approach to light curves for an X-ray binary, two AGN, a long-period variable star, and an RR-Lyrae star, in order to illustrate their use, applicability, and interpretation. ",
        "introduction": "\\label{s-intro}  Current and future time domain optical surveys, such as the SDSS Stripe 82 Supernova Survey \\citep{Frieman2008a}, Palomar Transient Factory \\citep[PTF,][]{Law2009a}, the Catalina Real-Time Transient Survey \\citep[CRTS,][]{Drake2009a}, Pan-STARRS \\citep{Kaiser2002a}, and the Large Synoptic Survey Telescope \\citep[LSST,][]{Ivezic2008a}, are providing and will provide an unprecedented flood of variability data. Such data sets will number in the millions to hundreds of billions of photometric data points, providing systematic multi-wavelength variability studies for thousands to eventually billions of objects. For example, LSST will have $\\sim 30$ trillion total measurements on $\\sim 40$ billion objects. For many classes of astronomical sources these will be the first systematic multi-passband variability studies including large numbers of objects with well-sampled multiwavelength light curves. Moreover, these rich data sets will also enable the inclusion of multi-passband variability information for distinguishing different classes of objects. Data sets generated by these surveys will present many exciting opportunities, providing astrophysical insight for known classes of objects, as well as the discovery of unknown variability classes, new subclasses of known variable classes, and anomalous outliers.  The central data analysis problem for extracting science from these time-domain data sets is how to quantitatively characterize the variability. For periodic signals characterizing the variability is relatively straightforward, with the period being the obvious and most important feature to use. For transient phenomenon, such as supernovae or black hole tidal disruptions, the light curve is often characterized by fitting a parameterized deterministic model to the data, either statistical (e.g., splines) or astrophysical. For quasi-periodic and stochastic light curves the variability is often characterized through the power spectral density (PSD). The PSD is the variability amplitude per frequency, so it describes the variability power contained within a frequency interval. A similar measure that is sometimes used is the structure function, which describes the variability amplitude as a function of time scale. The variability of quasi-periodic and stochastic light curves may then be characterized by summarizing the power spectrum through a parametric form, with power-laws and sums of Lorentzian functions being common choices for AGN and X-ray binaries, respectively.  Irregular sampling or sequences of regular sampling separated by gaps are often the source of the most problematic aspects of measuring variability features from a lightcurve. Unfortunately, all ground-based astronomical data and many space-based data are subject to sampling which is not strictly uniform. For periodic signals, methods have been developed to estimate periods from irregularly sampled light curves, and assess their statistical significance \\citep[e.g.,][]{Scargle1982a,Horne1986a,Reimann1994a}; significance tests have also been developed for periodic signals against red noise for regularly sampled light curves \\citep{Vaughan2010a}. For deterministic models, fitting the light curve is a traditional regression problem so the sampling pattern does not bias the results.  Deriving the PSD features and their uncertainties from an irregularly sampled light curve is considerably more challenging for stochastic light curves. For a regularly sampled light curve the traditional way of estimating the PSD is through the Discrete Fourier Transform of the light curve, the modulus-squared of which is called the periodogram. The periodogram suffers from biases due to the fact that it is calculated from a time series that is a sample from a continuous-time stochastic process. As a result, the sampling pattern of the light curve distorts the periodogram relative to the true PSD that generated the light curve \\citep[e.g.,][]{Uttley2002a, Vaughan2003a}. These issues similarly cause the empirical structure function to be a distorted estimate of the true structure function \\citep{Emmanoulopoulos2010a}. This is a concern because if one does not correct for this distortion, differences in variability properties caused by vagaries of the sampling pattern could mistakenly be interpreted to have an astrophysical origin. Moreover, this distortion is also a problem for classification algorithms that utilize variability information, as objects may spuriously be interpreted to belong to different variability classes simply because they have different sampling patterns. Clearly this issue must be dealt with in order to take advantage of the state of the art time domain data sets.  There are two primary approaches used in astronomical studies to account for irregular sampling in stochastic light curves. The first approach is to use Monte Carlo simulations to forward model the periodogram as a function of a model power spectrum \\citep{Done1992a,Uttley2002a,Emmanoulopoulos2013a}. The approach proceeds by first simulating a large number of light curves from an assumed power spectrum, sampling them with the same pattern as the measured light curve, computing the periodogram for each down sampled simulated light curve, and averaging the simulated periodograms to create a `model' periodogram. The best-fit power spectrum is found by minimizing the $\\chi^2$ between the model periodogram and the measured periodogram. Moreover, confidence intervals may be estimated from the Monte Carlo samples. This approach is extremely flexible, but can be computationally expensive. This is especially true if there are intervals of fine sampling separated by intervals of sparse sampling, as this requires either generating a very dense light curve at the finest sampling rate, or splitting the light curve into segments and computing their periodograms separately. Unfortunately, because of the computational cost, it is difficult to see how this approach can be applied to massive time domain data sets. Moreover, fitting complex multi-component for even a single data set can become computationally expensive in the Monte Carlo forward fitting approach.  The second approach to accounting for irregular sampling in stochastic light curves is to fit the light curve in the time domain. This is almost always done by assuming the light curve is a realization of a Gaussian process \\citep[e.g.,][]{Rybicki1992a,Kelly2009a,Kelly2011a,Kelly2013b,Miller2010a,Kozowski2010a,Brewer2011a}. In this case, the likelihood function is a multivariate normal distribution with unknown mean and covariance matrix. The covariance matrix is parameterized by the autocovariance function, which forms a Fourier transform pair with the power spectrum. In this time-domain approach, the autocovariance function is fit through maximum-likelihood or Bayesian approaches, and the power spectrum is calculated from the inferred autocovariance function. The irregular sampling and measurement errors are automatically accounted for by the likelihood function. This approach is statistically powerful because of its reliance on the likelihood function, but in general is also computationally expensive. In order to calculate the multivariate Gaussian likelihood function it is necessary to invert the $n \\times n$ covariance matrix of the light curve, where $n$ is the number of data points in the light curve. In general matrix inversion scales as $O(n^3)$, which can be prohibitive for large samples of light curves that are sampled with even moderate density.  There are special classes of Gaussian processes for which the computational complexity only scales linearly with the length of the light curve. For linear processes which have a `state-space' representation \\citep[discussed further in \\S~\\ref{s-statespace}, see also][]{Vio1992a} the likelihood function can be computed using a computationally efficient (scaling as $O(n)$) algorithm known as the Kalman Filter. Gaussian processes with an exponential autocorrelation fall into this class of stochastic models that have fast algorithms for computing their likelihood function \\citep{Rybicki1995a,Kozowski2010a}. This particular process is known as both a first order continuous-time autoregressive process (CAR(1)) or an Ornstein-Uhlenbeck process\\footnote{There is effectively no difference between an Ornstein-Uhlenbeck process and a first-order continuous time autoregressive process, although the terminology Ornstein-Uhlenbeck process is more common in the physics literature.}, and was introduced by \\cite{Kelly2009a} as a model for quasar optical light curves. Subsequent work has confirmed that it provides a good description of quasar optical light curves on time scales of days to years at the level of data quality of the OGLE and Stripe 82 surveys \\citep{Kozowski2010a,MacLeod2010a,Zu2013a,Andrae2013a}. However, recent studies have found evidence for deviations from the CAR(1) model for optical light curves of AGN \\citep{Mushotzky2011a,Graham2014a}. In addition, successful AGN selection techniques have been developed based on the CAR(1) parameters \\citep{Butler2011a,MacLeod2011a,Ruan2012a,Choi2013a}.  Despite its recent success, the CAR(1) model is very simple, as it assumes a PSD that is a Lorentzian centered at zero. Thus, there are only two free parameters in the CAR(1) model: the bend frequency of the PSD (i.e., the width of the Lorentzian) and the normalization. This makes the model inflexible, limiting its broader use. In this paper we overcome the inflexiblity of the CAR(1) model and present the general class of continuous time autoregressive moving average (CARMA) models. CARMA models are generated by adding higher order derivatives to the stochastic differential equation that defines the CAR(1) processes. The special case of a CAR($p$) process is discussed by \\citet{Koen2005a}, who applied it to the light curves for two variable stars. The PSD of a CARMA process is a sum of Lorentzian functions, with the free parameters being the centroid, widths, and normalizations of the Lorentzian functions. This provides a significant amount of flexibility in modeling the PSD, making CARMA modeling applicable to many classes of astronomical variables. Moreover, CARMA models have a state space representation, enabling the use of the Kalman Filter for calculating their likelihood function. Because of this, the computational complexity of calculating the likelihood function for CARMA models still scales linearly with the number of data points in the light curve, making them scalable to massive time domain data sets.  The format of this paper is as follows. In \\S~\\ref{s-carma_def} we begin by defining the CARMA process via a stochastic differential equation, present the PSD of a CARMA process, and present the autocovariance function of a CARMA process. Then, in \\S~\\ref{s-statespace} we express the CARMA process using a continuous-time state space representation and use this representation to derive the solution to the stochastic differential equation defining the process. This solution forms the basis for the statistical properties of a CARMA process sampled at a set of observational times, which are necessary for fitting the CARMA process parameters to a measured time series. In \\S~\\ref{s-inference} we present the likelihood function for a sampled CARMA process, and in \\S~\\ref{s-likelihood} present an efficient algorithm for computing the likelihood function based on the Kalman filter. In order to derive the Kalman filter for a CARMA process it is necessary to obtain a discrete-time state space representation of the sampled process, and we begin \\S~\\ref{s-likelihood} by using the results of \\S~\\ref{s-statespace} to present this discrete time representation. We extend these results in \\S~\\ref{s-interpolation} and derive an algorithm for efficiently performing interpolation and extrapolation from a measured time series assuming a CARMA model. In \\S~\\ref{s-bayes} we describe Bayesian inference for the CARMA model, including our adopted prior distribution, in \\S~\\ref{s-assess_fit} we describe how to assess the quality of fit based on the CARMA model, in \\S~\\ref{s-order} we discuss how to choose the order of the CARMA model, and in \\S~\\ref{s-computation} we discuss computational aspects related to fitting the CARMA model to a measured time series. In \\S~\\ref{s-simulations} we illustrate statistical inference under a CARMA model on two simulated lightcurves, and in \\S~\\ref{s-applications} we apply the CARMA model to astronomical lightcurves from an X-ray binary, AGN, and periodic variable stars. In \\S~\\ref{s-discussion} we discuss our results, and provide directions for future development. ",
        "conclusions": "\\label{s-discussion}  In this work we have introduced continuous-time autoregressive moving average processes as flexible models for stochastic and quasi-periodic light curves. These models account for irregular sampling and measurement errors, making them applicable to a wide variety of light curves. Moreover, they are flexible, as their PSD can be described as a sum of Lorentzian functions. The primary purposes of this modeling approach are 1) to provide a flexible way to estimate power spectra for astronomical light curves, and 2) to provide variability features for light curves that may be used in variability selection techniques and potentially in the identification of new classes of variables. Because one can compute the likelihood function for a light curve under a CARMA model they have the advantage that they are statistically efficient and rigorous, as all of the information in the light curve is used to estimate the variability parameters and inference is based on the well-developed statistical theory of maximum-likelihood or Bayesian inference. Moreover, calculation of the likelihood function is computationally efficient, scaling linearly with the number of data points in a light curve. This last point makes their application to massive time domain data sets particularly attractive.  Previous work has also expanded upon the CAR(1) model to introduce additional flexibility. \\citet{Kelly2011a} developed a mixture of CAR(1) processes as a model for X-ray variability of AGN. In this model, the light curve is expressed as a weighted sum of independent CAR(1) processes with different characteristics time scales that are constrained to lie on a regular logarithmic grid. The free parameters for this model are the maximum and minimum characteristic time scales of the grid, the mean and variance of the lightcurve, and the weights. Comparison with Equation (\\ref{eq-carma_autocovar} that the mixture of $p$ CAR(1) processes model of \\citet{Kelly2011a} is a special case of a CARMA(p,q) process where the roots of Equation (\\ref{eq-characteristic}) are constrained to be real. \\citet{Kelly2011a} also showed that a mixture of CAR(1) processes can closely approximate a broken-power law model for the PSD. In this case, the maximum and minimum characteristic time scales correspond to the low and high frequency breaks, respectively, and the the sequence of weights can be calculated as a function of the slope of the power-law between the low and high frequency breaks, so long as the slope of the PSD is constrained to the range $[-2, 0]$. These constraints reduce the number of free parameters to five for this model, which can be a computational advantage. In addition, \\citet{Kelly2011a} also showed that the solution to the stochastic linear diffusion equation is a mixture of CAR(1) processes, providing a physical interpretation of the variability model.  \\citet{Andrae2013a} investigated extensions to the Gaussian CAR(1) process as a model for quasar variability. The processes investigated by \\citet{Andrae2013a} included models with more flexible PSDs, such as a CAR(2) process and a CARMA(1,1) process. The former process provides the ability to capture quasi-periodic variations, while the latter process introduces additional smoothing of the stochastic driving noise. Both processes are special cases of the general class of CARMA processes we discuss here; however, we note that the CARMA(1,1) model is not stationary. \\citet{Andrae2013a} also investigated CAR(1) models after relaxing the assumptions of a linear Gaussian process. In particular, they also investigated non-Gaussian CAR(1) models and non-linear processes where the light curve variance also stochastically varies in time. These models, while not as flexible in modeling the PSD as a CARMA process, provide valuable alternatives to linear Gaussian models and may provide a better description of the variability of some light curves. However, in the case of quasar light curves from Stripe 82 \\citet{Andrae2013a} concluded that the linear Gaussian CAR(1) process provided the best model for most of the quasars, based on a Bayesian model comparison.  While CARMA models are, in theory, computationally efficient, the likelihood space can be complex and exhibit multiple modes which can be problematic for numerical optimizers and samplers. This is especially true for higher order models, and seems to affect higher order $q$ models more strongly than higher order $p$ models. The likelihood space may also be very complex for light curves that have regular nearly deterministic variations, such as RR-Lyrae stars. We consider the difficulty in optimizing or sampling from a multimodal complex posterior to be the primary short-coming of these models at this time, and dampens their computational efficiency. Thus, researchers who utilize them must be careful to check that the primary posterior mode has been found, and that the dominant posterior modes have been sampled from. We deal with this in our MCMC sampler using a parallel tempering algorithm, which we have used with success. However, even this can fail to adequately sample the posterior if an insufficient number of chains are run, or if the algorithm is not run for a sufficiently long period of time. Moreover, our maximum-likelihood estimation is rather simple, as we simply use 100 random starts and optimize by finding a local mode using a greedy gradient-based algorithm. Future work should focus on improving the effectiveness and efficiency of the maximum-likelihood estimate and MCMC algorithm.  In order to illustrate the applicability of CARMA models for a variety of astronomical sources, as well as to provide a guide for interpreting their results, we applied these models to light curves for an X-ray binary, two AGN, a long-period variable star, and a RR-Lyrae star. In general we found that the CARMA models provide a good description of these light curves, suggesting that they can be applied to a broad range of astronomical sources that exhibit stochastic or quasi-periodic variations. The only exception was the RR-Lyrae star light curve. This is to be expected, as RR-Lyrae light curves exhibit regular non-sinusoidal variations and thus are unlikely to have a strong stochastic component. However, in spite of this the CARMA models still identified the period quoted by the catalogue that this light curve was taken from in $\\approx 75\\%$ of the MCMC samples. Moreover, the deviation in the residuals from a normal distribution for the RR-Lyrae star implies that the distribution of the residuals from a CARMA fit may provide an effective means of discriminating between different types of variables, even for those for which a CARMA model is not optimal. These results suggests that there may be value in using variability features derived from CARMA parameters even for non-stochastic light curves.  Further improvements to the CARMA modeling approach can be obtained by including a deterministic component, such as a periodic function, for modeling periodic light curves, such as those from RR-Lyrae. In this case the residuals from fitting a deterministic function are modeled as following a CARMA process. In fact, this is the motivation behind the Periodic Autoregressive Moving Average models \\citep[PARMA, e.g.,][]{Anderson2013a}, which allow for periodic variations in the mean and autocovariance function of a time series. In addition, it is possible to define multivariate CARMA models through a vector- and matrix-valued extension to Equation (\\ref{eq-carma}) \\citep[e.g.,][]{Marquardt2007a,Schlemm2012a}. Multivariate CARMA models hold considerable potential for characterizing the full multi-passband variability information obtained by time-domain surveys, and will be the subject of future work.  In summary, CARMA models provide an important addition to the astronomer's statistical toolbox in the era of massive time-domain surveys, and have the potential to play an important role in the analysis of variability as a probe of astrophysics, as well as in the use of variability as a means of identifying classes of astronomical sources."
    },
    "1402/1402.2769_arXiv.txt": {
        "abstract": "In a uniformly rotating fluid, inertial waves propagate along rays that are inclined to the rotation axis by an angle that depends on the wave frequency. In closed domains, multiple reflections from the boundaries may cause inertial waves to focus on to particular structures known as wave attractors. These attractors are likely to appear in fluid containers with at least one boundary that is neither parallel nor normal to the rotation axis. A closely related process also applies to internal gravity waves in a stably stratified fluid. Such structures have previously been studied from a theoretical point of view, in laboratory experiments, in linear numerical calculations and in some recent numerical simulations. In the present paper, two-dimensional direct numerical simulations of an inertial wave attractor are presented. By varying the amplitude at which the system is forced periodically, we are able to describe the transition between the linear and nonlinear regimes as well as the characteristic properties of the two situations. In the linear regime, we first recover the results of the linear calculations and asymptotic theory of \\cite{Ogilvie05} who considered a prototypical problem involving the focusing of linear internal waves into a narrow beam centred on a wave attractor in a steady state. The velocity profile of the beam and its scalings with the Ekman number, as well as the asymptotic value of the dissipation rate, are found to be in agreement with the linear theory. We also find that, as the beam builds up around the wave attractor, the power input by the applied force reaches its limiting value more rapidly than the dissipation rate, which saturates only when the beam has reached its final thickness. In the nonlinear regime, the beam is strongly affected and becomes unstable to a subharmonic instability. This instability transfers energy to secondary waves possessing shorter wavelengths and lower frequencies. The onset of the instability of a narrow inertial wave beam is investigated by means of a separate linear analysis and the results, such as the onset of the instability, are found to be consistent with the global simulations of the wave attractor. The excitation of such secondary waves described theoretically in this work has also been seen in recent laboratory experiments on internal gravity waves. ",
        "introduction": "Internal waves supported by buoyancy and Coriolis forces play an important role in the Earth's ocean and atmosphere, and have been widely studied through theory and experiment.  The simplest situations are those of pure internal gravity waves in a non-rotating, stably stratified fluid, and pure inertial waves in a rotating, unstratified fluid.  There is a close analogy between these two cases.  Internal waves have properties quite different from acoustic or electromagnetic waves, being highly anisotropic and dispersive.  In a frame of reference in which the mean flow vanishes, the wave frequency depends only on the direction, and not on the magnitude, of the wavevector.  Special directions are defined by gravity or rotation. Energy propagates at the group velocity, which is perpendicular to the wavevector and inversely proportional to its magnitude. Internal waves are highly dispersive and do not steepen like acoustic waves, but instead break and dissipate when they exceed a critical amplitude and become locally unstable. Their frequencies are bounded and form a dense or continuous spectrum.  In a closed domain, internal waves generally do not form normal modes with regular eigenfunctions, although the simplest cases are exceptions to this rule.  Waves of a given frequency generally undergo focusing or defocusing when they reflect from boundaries; after multiple reflections in a generic container, rays are typically focused towards limit cycles known as wave attractors \\citep{Maas95}.  In the interiors of stars and giant planets, internal waves occupy the low-frequency part of the spectrum of oscillations, which can be excited by tidal forcing when the body has an orbiting companion \\citep[e.g][]{Ogilvie04}.  The question arises of how a rotating or stably stratified fluid responds to an oscillatory body force with a frequency that lies within the dense or continuous spectrum.  For tidal applications, what is of particular interest is the rate at which energy is transferred to the fluid, ultimately to be dissipated.  The relative amplitudes of astrophysical tides can vary enormously, depending on the relative mass and distance of the tide-raising body.  It is therefore appropriate to investigate both linear and nonlinear regimes for internal tides.  \\cite{Ogilvie05} considered a simple problem involving a two-dimensional fluid domain that supports internal waves, and calculated the steady-state linear response to a periodic body force with a frequency such that a wave attractor occurs.  He obtained an analytical asymptotic solution in the limit that the dissipative effects of the fluid are small. The response is localized in a narrow beam centred on the attractor, and, for a body force of fixed amplitude, the dissipation rate (which is equal to the power input) is asymptotically independent of both the magnitude and the details of the dissipative process. This situation occurs because the attractor acts as a `black hole' that absorbs a certain energy flux from the waves being focused into it. This result contrasts strongly with a situation in which the internal waves form classical normal modes, as in the case of inertial waves in a rectangular container whose walls are parallel and perpendicular to the rotation axis. In such a situation, the dissipation rate is simply proportional to the viscosity, and therefore astrophysically unimportant, unless the forcing frequency matches the natural frequency of a normal mode and a resonance occurs.  Internal wave attractors have been seen in experiments \\citep{Maas97,Maas01,Manders03, Lam08, Hazewinkel08, Hazewinkel11a, Hazewinkel11b,Scolan13}, linear calculations \\citep{Rieutord97, Dintrans99, Rieutord01, Ogilvie05, Maas07, Ogilvie09, Rieutord10, Echeverri11, Bajars13, Baruteau13} and numerical simulations \\citep{Grisouard08}.  The wave attractor is a linear phenomenon and one may naturally ask how it is modified by nonlinearity.  Previous experiments \\citep{Maas01} have shown some evidence for cumulative nonlinear effects that build up over a large number of wave periods: the fluid becomes increasingly mixed, and momentum deposition leads to the generation of a mean flow.  Until recently, no evidence had been found for stronger nonlinear phenomena such as wave breaking, instability or turbulence.  However, \\citet{Scolan13} have recently reported experiments on internal gravity waves in which waves focused towards an attractor become unstable and divert energy to smaller scales.  They describe this process as a parametric subharmonic instability similar to that known for plane internal waves.  The purpose of this work is to study a simple system of periodically forced inertial waves in a closed domain by means of fully nonlinear direct numerical simulations.  We aim first to recover the results of \\cite{Ogilvie05} for the steady-state response of a wave attractor in a linear regime, and to understand how this steady state is achieved. Moving into the nonlinear regime, we aim to locate the onset of instability and to see how the attractor operates, if at all, in the presence of instability.  We will compare our results with the experimental findings of \\citet{Scolan13}.  Of particular interest is the energy budget of the system and the question of how the dissipation rate (and thence the tidal torque, in astrophysical applications) is affected by nonlinearity. ",
        "conclusions": "The purpose of this work was to simulate numerically the evolution in 2D of an inertial wave attractor in a square container in the linear and nonlinear regimes, with a particular focus on the build-up of the attractor before it reaches a steady-state and the instability occurring in the nonlinear case. This study complements previous work on this problem, namely laboratory experiments and linear calculations. In the linear regime, the characteristics of internal wave attractors predicted by theory are well recovered. In particular, we find that the thickness of the attractor scales with the Ekman number to the power $1/3$. For a fixed Ekman number, we moreover confirm that after a reflection at the boundary, the amplitude of the velocity field decreases as the distance from the source to the power $1/3$ while the breadth of the wave beam increases by the same factor. Finally, we find in the simulations that, for a body force of fixed amplitude, the dissipation rate indeed becomes independent of the Ekman number and approaches the asymptotic value derived analytically in \\cite{Ogilvie05}.  Since we compute the evolution of the wave attractor from the very beginning where an incompressible fluid is forced in a closed container, we have access to the transient phase before the system reaches a statistically steady state. In this simple linear case, only two quantities will play a role in the energy balance: the power input due to the periodic forcing of the system and the dissipation rate. Both those quantities reach the same absolute value in the steady state in order for energy to be conserved. However, we find here that the power (i.e. the rate of working by the applied force) saturates in about 5 wave periods in all cases, while the dissipation rate reaches its asymptotic value on a timescale that depends on the Ekman number, typically close to 15 wave periods. The kinetic energy of the flow therefore continues to build up, and the dissipation rate increases, until the wave attractor has reached its final thickness, while the rate of working by the applied force saturates much earlier. This is consistent with the viewpoint that the attractor acts as a ``black hole'' that absorbs the energy flux associated with the quasi-inviscid, large-scale waves that are focused towards it.  The present work also focused on the properties of inertial wave attractors in the nonlinear regimes. To do so, the forcing amplitude was increased until nonlinear effects would start to play a significant role. The primary wave beam oscillating at the forcing frequency $\\omega_0$ is then found to be subject to an instability which is investigated through linear analysis and nonlinear calculations of the simpler case of a single infinite narrow wave beam. The onset of this instability and the nonlinear outcome are determined in this simpler case and are consistent with the characteristics of the global wave attractor in the nonlinear regime. It is found that this instability transfers energy to waves of lower frequencies and shorter wavelengths. Those secondary inertial waves are excited close to the reflection points and then propagate with a direction in agreement with the typical inertial wave dispersion relation. Since they possess a smaller scale, they dissipate quickly and propagate over much shorter distances than the primary waves. As the forcing amplitude is increased, their effective Reynolds number is increased and the frequencies of the secondary waves drift towards $\\omega_0/2$, thus approaching the outcome of a regular parametric subharmonic instability in the inviscid case. In the cases where instability occurs, the energy is transferred rapidly to shorter wavelengths which then dissipate. As a consequence, the dissipation rate reaches its saturation more quickly than in the linear cases. In the present work, this type of instability is described theoretically and recovered in nonlinear direct numerical simulations. An instability of a wave beam with similar characteristics was found in an experiment and supported by nonlinear numerical simulations by \\cite{Clark10}. Other types of instabilities were also observed in laboratory experiments of gravity waves \\citep{Joubaud12} or inertial waves \\citep{Bordes12} excited in closed containers.  In geophysical or astrophysical situations, inertial waves can be excited in stars or planets by tidal forcing due to an orbiting companion. Moreover, these systems can easily be in a nonlinear regime if the forcing is sufficiently strong (because of the companion being sufficiently close and/or massive). Under these conditions, wave attractors can still be present but the flow may be unstable and the dissipation rate and tidal torque may be reduced below the value predicted by linear theory.  Of course the problem considered here was 2D, consisting of a very simple geometry and it has been shown that moving to 3D and spherical geometry may strongly modify the behaviour of wave attractors and the characteristics of the dissipation processes \\citep{Rieutord10}. However, it is the first example here of strong nonlinear effects emerging in simulations of wave attractors, namely subharmonic instabilities. It is interesting to note that our simulations have counterparts in the recent fluid experiments of \\cite{Scolan13}, which also exhibit a strongly nonlinear regime of wave attractors.  Obviously, much more can be done on this problem. If we consider the conditions in the interiors of stars or giant planets, large-scale flows such as meridional flows or differential rotation are likely to exist and modify the build-up of potential wave attractors. Turbulent convection will of course also probably act on the wave focusing and defocusing and may disrupt the formation of wave attractors. A next possible step could be the computation of the same kind of simulations for a compressible fluid in a rotating spherical shell, where both gravity and inertial waves would exist. In spherical geometry, critical latitude singularities were shown to be dominant features of periodically forced flows in a linear regime. It would be valuable to carry out fully nonlinear simulations of these phenomena."
    },
    "1402/1402.6319_arXiv.txt": {
        "abstract": "We report the discovery of an object near M87 in the Virgo Cluster with an extraordinary blueshift of $-1025$ \\kms, offset from the systemic velocity by $>2300$ \\kms. Evaluation of photometric and spectroscopic data provides strong evidence that this object is a distant massive globular cluster, which we call HVGC-1 in analogy to Galactic hypervelocity stars. We consider but disfavor more exotic interpretations, such as a system of stars bound to a recoiling black hole. The odds of observing an outlier as extreme as HVGC-1 in a virialized distribution of intracluster objects are small; it appears more likely that the cluster was (or is being) ejected from Virgo following a  three-body interaction. The nature of the interaction is unclear, and could involve either a subhalo or a binary supermassive black hole at the center of M87. ",
        "introduction": "Extreme astrophysical objects are occasionally found in large samples of data. While classifying a million galaxy images, the Galaxy Zoo project found Hanny's Voorwerp, an unusual cloud ionized by an active galactic nucleus \\citep{Lintott}. In a study of blue horizontal branch stars in the Galactic halo selected from the Sloan Digital Sky Survey, \\cite{brown2005} found hypervelocity stars that are thought to originate in three-body interactions with the supermassive black hole (SMBH) at the center of the Galaxy. In an extensive study of star clusters, planetary nebulae, and \\ion{H}{2} regions in M31, \\cite{caldwell10} found one star with the most negative velocity known ($-780$~\\kms), a probable member of the Andromeda giant stream.  Here, we report on an object found in a different large survey that has an even more extreme negative velocity: an apparent globular cluster (GC) toward the central Virgo Cluster galaxy M87. ",
        "conclusions": ""
    },
    "1402/1402.1680_arXiv.txt": {
        "abstract": "The Schmidt-Kennicutt relation links the surface densities of gas to the star formation rate in galaxies. The physical origin of this relation, and in particular its break, i.e. the transition between an inefficient regime at low gas surface densities and a main regime at higher densities, remains debated. Here, we study the physical origin of the star formation relations and breaks in several low-redshift galaxies, from dwarf irregulars to massive spirals. We use numerical simulations representative of the Milky Way, the Large and the Small Magellanic Clouds with parsec up to subparsec resolution, and which reproduce the observed star formation relations and the relative variations of the star formation thresholds. We analyze the role of interstellar turbulence, gas cooling, and geometry in drawing these relations, at 100 pc scale. We suggest in particular that the existence of a break in the Schmidt-Kennicutt relation could be linked to the transition from subsonic to supersonic turbulence and is independent of self-shielding effects. This transition being connected to the gas thermal properties and thus to the metallicity, the break is shifted toward high surface densities in metal-poor galaxies, as observed in dwarf galaxies. Our results suggest that together with the collapse of clouds under self-gravity, turbulence (injected at galactic scale) can induce the compression of gas and regulate star formation. ",
        "introduction": "Star formation is among the most important physical processes affecting the formation and evolution of galaxies. Nevertheless, fundamental questions about the efficiency of the conversion of gas into stars and what triggers the process of star formation remain open.  Observations of galaxies have shown a close correlation, known as the Schmidt-Kennicutt relation, between the surface density of star formation rate (\\SigmaSFR) and the surface density of gas (\\Sigmagas) \\citep[e.g.][]{Kennicutt1989, WongBlitz2002}. The details of this scaling relation are found to vary with the environment. Spiral galaxies convert their gas into stars with longer depletion times than galaxies in a merger phase \\citep{Daddi2010, Genzel2010, Saintonge2012}, but more rapidly than dwarf galaxies \\citep{Leroy2008,Bolatto2011}. In addition, the observed relation varies for different tracers. It is shallower for molecular gas than for total (molecular and atomic) gas \\citep{GaoSolomon2004, Bigiel2008, Heyer2004}, but steeper when the atomic gas is considered solely \\citep{Kennicutt1998, Kennicutt2007, Schuster2007, Bigiel2008}. Several other observations \\citep[e.g.][]{Kennicutt1989, MartinKennicutt2001, Boissier2003} have suggested the existence of a critical surface density, the so-called break, below which \\SigmaSFR{} in spiral galaxies drops: star formation is inefficient compared to the regime at high surface densities, well described by a power-law. Consequently, a composite relation of star formation seems to be a better description for the \\Sigmagas--\\SigmaSFR{} relation, rather than a single power-law.  However, the origin of the break, and the transition to a different regime of star formation at high surface densities, remain a matter of debate. Some models evoke the Toomre criterion for gravitational instability \\citep[e.g.][]{Quirk1972, Kennicutt1989, MartinKennicutt2001} or rotational shear \\citep{Hunter1998,MartinKennicutt2001} to interpret the existence of the break. \\cite{Elmegreen1994} emphasize the need for the coexistence of two thermal phases in pressure equilibrium and \\cite{Schaye2004} argues that it is the transition from the warm to the cold gas phase, enhanced by the ability of the gas to shield itself from external photo-dissociation, that triggers gravitational instabilities over a wide range of scales. Self-shielding plays an important role also in the model of \\cite{Krumholz2009}, where it sets the transition from atomic to molecular phase at a metallicity-dependent \\Sigmagas. \\cite{DibPMB2011} shows that feedback from massive stars is an important regulator of the star formation efficiency in protocluster forming clouds. Based on this, \\cite{Dib2011} proposes that the break in the \\Sigmagas--\\SigmaSFR{} relation can be related to a feedback-dependent transition of the star formation efficiency per unit time, as a function of the gas surface density (but see \\citealt{Dale2013} reporting a possibly low impact of the stellar feedback on the star formation rate and efficiency). \\cite{Renaud2012} have recently proposed an analytic model in which the origin of the star formation break is related to the turbulent structure of the interstellar medium (ISM). In this model, a threshold in the local volume density, resulting in the observed surface density break corresponds to the onset of supersonic turbulence that generates shocks which in turn trigger the gravitational instabilities and thus star formation.  Different mechanisms are invoked in theoretical works to explain the scaling relations, such as gravity \\citep{Tan2000, SilkNorman2009}, turbulence of the interstellar medium \\citep[e.g.][]{Elmegreen2002, MacLowKlessen2004, Krumholz2005, Hennebelle2011, Padoan2011, Federrath2012, Renaud2012, Federrath2013}, feedback from massive stars \\citep{Dib2011} and the interplay between the dynamical and thermal state of the gas \\citep{Struck1999}.  In addition to these theoretical studies, several galaxy simulations modeling the ISM found a reasonable agreement with observations, using various recipes for star formation and stellar feedback \\citep{Li2006, WadaNorman2007, RobertsonKravtsov2008, TaskerBryan2008, DobbsPringle2009, KoyamaOstriker2009, Agertz2011, Dobbs2011, Kim2011, Monaco2012, Rahimi2012, ShettyOstriker2012,Halle2013}. Among them \\cite{Bonnell2013} resolved the small scale physics of star formation in the context of galactic scale dynamics. The observed correlation between \\Sigmagas{} and \\SigmaSFR, together with the break of \\SigmaSFR{} are often reproduced in simulations, but it remains unclear to what extent the star formation rate estimates depend on the parameters of individual models and underlying assumptions, and what are the fundamental drivers for the observed relations.  In this paper we aim at providing a better understanding of the star formation relations and threshold by studying the local properties of simulated galaxies. Our work is in great part motivated by the analytic model of \\cite{Renaud2012}, based on the supersonic nature of the turbulence in the ISM. Their formalism leads to an analytic expression relating \\Sigmagas{} and \\SigmaSFR{} that depends on three parameters: the Mach number, the star formation density threshold and the thickness of the star-forming regions. One assumption of this model is the characterization of an entire star-forming region by a single set of these three parameters, while wide ranges of them are more appropriate for describing the real ISM. Another assumption is the description of the gas volume density by a log-normal distribution, which was primarily found for isothermal supersonic turbulence \\citep[e.g.][]{Vazquez1994,NordlundPadoan1999}. By performing galaxy simulations, we achieve a wide diversity of parameters and density distributions consistent with the multiphase ISM. Simulations allow us to study local properties of individual regions of galaxies, such as velocity dispersion, temperature or geometry and to infer their impact on the \\Sigmagas{} -- \\SigmaSFR{} relation.  We start with the presentation of the simulation technique and details of galaxy models in Section \\ref{Sec:Simulations}. Section \\ref{Sec:3galaxies} presents the analyzed sample of galaxies. The method used in deriving parameters needed for the analysis is described in Section \\ref{Sec:Analysis}. The dependence of star formation on different parameters, plotted in the \\Sigmagas{} -- \\SigmaSFR{} parameter space, is shown in Section \\ref{Sec:Results}. In Section \\ref{Sec:Discussion}, we discuss the results and compare with the model of \\cite{Renaud2012}. Finally, we conclude with the summary in Section \\ref{Sec:Summary}. ",
        "conclusions": "\\label{Sec:Discussion}  \\subsection{Comparison with observations} \\label{Sec:Comparison}  Most spatially resolved studies of spiral galaxies find the presence of a power-law \\Sigmagas--\\SigmaSFR{} relation with a break at surface gas densities of the order of a few M$_{\\odot} \\, \\mathrm{pc}^{-2}$ \\citep[see][and references therein]{KennicuttEvans2012}. The slope of the power-law relation in the high surface-density regime is found to be in the range 1.2--1.6 when total (molecular plus atomic) gas surface density is considered.  Less agreement about the power-law slope in observations is reached when molecular-gas surface density is considered solely. Some recent studies \\citep[e.g.][]{Eales2010,Rahman2011,Leroy2013} have reported an approximately linear relation between the surface density of star formation rate and the surface density of molecular gas. Other studies \\citep[e.g.][]{Kennicutt2007,Verley2010,Liu2011} have found a much steeper relation, with a slope in the range 1.2--1.7, similar to that of integrated measurements \\citep{Kennicutt1998}. This discrepancy between different results in observations is still debated. A possible interpretation of the sublinear relation was recently proposed by \\cite{Shetty2013}. They suggest that the CO emission used in the estimation of \\Sigmagas{} is not all necessarily associated with SF. Not subtracting off such a diffuse component could lead to a slope close to unity.  The distribution of data points from the observations of the \\smc{} \\citep{Bolatto2011} in the \\Sigmagas--\\SigmaSFR{} plane has a similar shape than that of spiral galaxies, but is noticeably shifted toward higher total \\Sigmagas.  In Figure \\ref{Fig:contour_all}, we show three of our models: \\mwpc, \\smcZth{} and \\lmcZsun{} in the \\Sigmagas--\\SigmaSFR{} plane. The \\mwpc{} and the \\lmcZsun{} models lie in the loci of observed spiral galaxies \\citep[e.g.][]{Kennicutt1998,Kennicutt2007,Bigiel2008}. Our \\smcZth{} model has a lower \\SigmaSFR{} for a given \\Sigmagas{} when compared to both the \\mwpc{} and the \\lmcZsun{} models. The region below the break of our \\smcZth{} model is located at slightly lower \\Sigmagas{} than the real Small Magellanic Cloud, but its displacement with respect to spiral galaxies (\\mw{} and \\lmc) is well reproduced (Figure \\ref{Fig:contour_all}). In our simulations, Equation \\ref{Eq:SFR} sets the slope of power-law relation with the value of 1.5. A shallower relation, closer to the observed values, might be reached by accounting for a stronger regulation of star formation (e.g. pre-SN stellar feedback, see \\citealt{Renaud2012}), but this slope change has not been demonstrated by simulations yet.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{contour_all3_3} \\caption{\\label{Fig:contour_all} Comparison of \\mwpc, \\lmcZsun{} and \\smcZth. The color coding of contour levels is as in Figure \\ref{Fig:MW_beam_compare_all}.} \\end{center} \\end{figure}   \\subsection{Interpretation of threshold} The \\textit{existence} of the break in the \\Sigmagas--\\SigmaSFR{} relation is, in our models, equivalent of having a non-zero value of the volume density threshold in the local, three dimensional star formation relation. Setting no threshold leads to a power-law relation without a break.  Figure \\ref{Fig:th_LMC} shows that the value of the density threshold $\\rho_0$ that we have used in our analysis has an impact on the \\Sigmagas--\\SigmaSFR{} relation. Changing the value of $\\rho_0$ changes the slope at low \\Sigmagas{} in the \\Sigmagas--\\SigmaSFR{} relation (Figure \\ref{Fig:th_LMC}). This could suggest that the transition from the inefficient to the power-law regime could be due to the density threshold $\\rho_0$ we imposed by hand in the star formation law (see Equation \\ref{Eq:SFR}). However, we have checked that the deviation from the power-law regime occurs at $\\mathcal{M} \\approx$ 1, independently of $\\rho_0$. In addition, in Figure \\ref{Fig:MW_SMC_LMC_Mach}, we have shown that beams located in the break tend to have $\\mathcal{M}$ below unity, while regions at high \\Sigmagas{} are mostly supersonic.  \\begin{figure}[t!] \\begin{center} \\includegraphics[width=\\columnwidth]{compare_th5_2} \\caption{\\label{Fig:th_LMC} Comparison of the \\Sigmagas--\\SigmaSFR{} relation in the \\lmcZsun{} simulation with the star formation volume density threshold $\\rho_0$=10 \\Hcc{} (black contours) to that with $\\rho_0$=100 \\Hcc{} (colored filled contours). The color coding of contour levels is as in Figure \\ref{Fig:MW_beam_compare_all}.} \\end{center} \\end{figure}  To better understand the behavior of the ISM in our simulations, we show in Figure \\ref{Fig:Mach_rho_T} the Mach number as a function of average volume density of gas\\footnote{computed as the mass-weighted average density of the gas in each beam} $\\left<\\rho\\right>$ in the beam for \\mwpc. The Mach number varies with the average density as $\\mathcal{M} \\propto \\left<\\rho\\right>^{0.5}$, similarly to the two-phase turbulent flow studied by \\cite{AuditHenne2010}. Although caution should be used when doing such comparison (temperature and velocity dispersion vary with density differently in both models), the onset of the supersonic regime, i.e. the transition from an inefficient regime to a power-law, happens at densities of $\\approx 10$ \\Hcc \\cite[see also][their Figures 4 and 9]{AuditHenne2010}.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{MW_Mach_rho} \\caption{\\label{Fig:Mach_rho_T} The Mach number as a function of the average volume density of gas computed in beams of 100$\\times$100$\\times$100 pc$^3$, for \\mwpc. The solid line corresponds to $\\mathcal{M} \\propto \\left<\\rho\\right>^{0.5}$ \\cite[][Figure 9]{AuditHenne2010}.} \\end{center} \\end{figure}  Other interpretations of the observed break are possible. The \\Sigmagas--\\SigmaSFR{} relation could be an effect of the galactic radial distance with low \\Sigmagas{} at large radius and high \\Sigmagas{} at small radius, as found by \\cite{Kennicutt2007} and \\cite{Bigiel2008}. The break could then be explained as a consequence of the drop in the average volume density in the outer regions of galaxies as proposed by \\cite{Barnes2012}. However, Figure \\ref{Fig:LMC_KS_radius2} shows no such radial dependence for \\Sigmagas{} nor \\SigmaSFR. Star-forming regions in outer parts of a galaxy can exhibit both star formation regimes. A possible explanation why we do not see any radial dependence in our simulations may be a missing metallicity gradient. The outer regions of our simulated galaxies have the same metallicity than the innermost regions, thus the metallicity is probably too high at the edge of disk and allows for an efficient cooling and consequently an efficient star formation while it may lie in the break regime otherwise.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{LMC_KS_radius2} \\caption{\\label{Fig:LMC_KS_radius2} The local surface density of the star formation rate as a function of surface density of gas for \\lmcZsun. Color indicates the radial distance of the beam in kpc.} \\end{center} \\end{figure}  When azimuthally averaged, \\Sigmagas{} and \\SigmaSFR{} both decline steadily as a function of radius in many galaxies despite different morphologies (see \\citealt{Bigiel2008} for a sample of nearby spiral galaxies and  \\citealt{Leroy2009} for CO intensity radial profiles for the same sample). Figure \\ref{Fig:LMC_profil} shows \\Sigmagas{} and \\SigmaSFR{} as functions of radius for \\lmcZsun{}. Both radial profiles decline with galactic radius as in observed spiral galaxies.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{LMC_profile} \\caption{\\label{Fig:LMC_profil} Radial profiles of azimuthally averaged \\Sigmagas{} (black) and \\SigmaSFR{} (red) for \\lmcZsun{}. Dotted lines correspond to exponential fits.} \\end{center} \\end{figure}  Another alternative explanation for the existence of the break is that it corresponds to the transition from atomic to molecular hydrogen \\citep{Krumholz2009}. According to this scenario, the transition from atomic to molecular hydrogen and the subsequent star formation depend on local conditions that vary with galactic radius, e.g. metallicity, gas pressure and shear\\footnote{Shear dictates whether molecular clouds can form \\citep[e.g.][]{Leroy2008,Elson2012}, but if they do form, shear does not seem to influence the efficiency at which they convert their gas into stars \\citep{Dib2012}.}. \\cite{Bigiel2008} found such radial dependence in the sample of nearby galaxies in agreement with the findings of \\citet{WongBlitz2002} and the threshold interpretations of e.g. \\citet{Kennicutt1989}, \\citet{MartinKennicutt2001} and \\citet{Leroy2008}. Similar results are reproduced in some simulations, e.g. \\cite{Halle2013}, who find that molecular gas is a better tracer of star formation than atomic gas and plays an important role in the low surface density regions of galaxies by allowing for more efficient star formation. However, our models that do not include chemodynamics, are able to reproduce the observed break at low $\\Sigmagas$. Therefore, this seems to indicate that the presence of molecules is not a necessary condition to \\textit{trigger} the process of star formation. However, we acknowledge numerous observational evidences showing that molecules are involved at a later stage of the SF process.  To summarize, we consider two representative beams having the same \\Sigmagas, but different \\SigmaSFR{} (Figure \\ref{Fig:LMC_rho_rho0_Mach}). These beams have similar average volume densities $\\left<\\rho\\right>$ which can be several orders of magnitude smaller than $\\rho_0$. However, the beam that happens to have the highest \\SigmaSFR{} has always the highest Mach number, as previously suggested by Figure \\ref{Fig:MW_SMC_LMC_Mach}. We have argued above that the density threshold $\\rho_0$, the thickness of the star-forming regions and the molecules do not have impact on the transition from the regime of inefficient star formation to the efficient power-law regime. The role of the artificial threshold $\\rho_0$ imposed in the simulations is to set a frontier between the diffuse non-star-forming gas and the star-forming component, but not to tune the efficiency of star formation per se. Therefore at a given \\Sigmagas, this efficiency depends mostly on the level of turbulence ($\\mathcal{M}$), i.e. the compression of the ISM.  \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=\\textwidth]{LMC_rho_rho0_Mach} \\caption{\\label{Fig:LMC_rho_rho0_Mach} The volume density normalized to the threshold density $\\rho_0$=10 \\Hcc{} (left panel) and the Mach number (right panel) for the \\lmcZsun{} simulation, as in Figure \\ref{Fig:MW_SMC_LMC_Mach}. Two pairs of beams (A$_1$--A$_2$ and B$_1$--B$_2$) are highlighted in each panel. Within each pair, the two beams are chosen to have the same value of \\Sigmagas{} and a similar value of $\\left<\\rho \\right>/ \\rho_0$. Beams with higher \\SigmaSFR{} at fixed \\Sigmagas{} have a higher Mach number. The two pairs of beams have larger size for clarity reasons.} \\end{center} \\end{figure*}  \\cite{Renaud2012} proposed that the break indeed corresponds to the onset of supersonic turbulence which, by generating shocks, triggers gravitational instabilities leading to star formation. In Figures \\ref{Fig:SMC_model} and \\ref{Fig:MW_model}, we compare simulations of \\mwpc{} and \\smcZth{} with the analytical model of \\cite{Renaud2012}. In this model, the relation between \\Sigmagas{} and \\SigmaSFR{} depends on three parameters: the Mach number $\\mathcal{M}$, the scale-height $h$ and the star formation volume density threshold $\\rho_0$. We do not compare each star-forming region in the simulation with the model, but we are rather interested in what values these parameters should take to obtain upper and lower limits for simulated data. The break is in the subsonic regime (measured values of $\\mathcal{M}$ are below unity) which corresponds to the regime where the analytical model deviates from its asymptotic behavior (at high \\Sigmagas). In this regime the scale-heights of the beams set the efficiency of star formation spanning the range given by the model and quantitatively in accordance with the values measured in the simulations (Figure \\ref{Fig:SMC_h}). In the analytical model, the power-law regime can be reached even with the Mach number below unity (red curve). However, our simulations do not probe this area of the \\Sigmagas--\\SigmaSFR{} plane: the data points in the power-law regime are exclusively supersonic and can only be described by a model with the Mach number above unity (black curve).  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{SMC_model_compare_new} \\caption{\\label{Fig:SMC_model} \\smcZth{}: comparison with the \\cite{Renaud2012} model with three sets of parameters (indicated in the legend). The black curve matches the supersonic regime of efficient star formation, while the green and the red curves represent upper and lower limits for the regime of the break, where the star formation is inefficient. The $h$ parameter is in agreement with values measured in the simulation (see Figure \\ref{Fig:SMC_h}).} \\end{center} \\end{figure}  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=\\columnwidth]{MW_model_compare2} \\caption{\\label{Fig:MW_model} \\mwpc{}: comparison with the \\cite{Renaud2012} model. As in Figure \\ref{Fig:SMC_model}, the supersonic regime is compared to the model prediction and similarly, the subsonic regime at low \\Sigmagas{} is situated between the curves characterized by the Mach number lower than unity for the measured thickness.} \\end{center} \\end{figure}  \\subsection{Metallicity}  In Figure \\ref{Fig:SMC_metallicity}, we have shown that the exact position of the break in the \\Sigmagas--\\SigmaSFR{} plane depends on metallicity. A comparison of different metallicities in otherwise identical systems shows that the slope at low \\Sigmagas{} has a greater value in metal-poor galaxies. Figure \\ref{Fig:3pdf} suggests that metallicity is not the only factor determining the gas density distribution in our simulations. Similar lack of direct dependence of the fraction of dense gas on metallicity is found when simulations of the \\smc{} with different metallicities are compared (not shown here). Thus the slope at low \\Sigmagas{} cannot be explained by the presence of a higher fraction of dense gas in systems with higher metallicity compared to systems with lower metallicity. However, metallicity has an impact on star formation, even though indirect. Metallicity directly influences the temperature of the gas: higher the metallicity, more efficient the cooling therefore the temperature and, in turn, impacts the Mach number. In Figure \\ref{Fig:SMC_metallicity2}, we show Mach numbers for \\smcZsun{} and \\smcZth{}. Higher values of Mach number are reached in galaxy with higher metallicity.  \\begin{figure*}[t!] \\begin{center} \\includegraphics[width=\\textwidth]{SMCs_compare_metallicity2} \\caption{\\label{Fig:SMC_metallicity2} As in Figure \\ref{Fig:SMC_metallicity}, the effect of gas metallicity on the \\Sigmagas--\\SigmaSFR{} relation in the model of \\smc{} is represented. Colors indicate the Mach number in the simulation of the \\smc{} with metallicity of 1.0 \\Zsun{} on the left panel and the simulation of the \\smc{} with metallicity of 0.1 \\Zsun{} on the right panel. In both panels the black contours are those of the simulation of the \\smcZth, shown for reference. } \\end{center} \\end{figure*}  This work does not include metallicity-dependent self-shielding and feedback. Accounting for them, \\cite{Dib2011} showed that both the fraction of gas in molecular form and the efficiency of star formation per unit time depend on metallicity. This leads to the metallicity dependent \\Sigmagas--\\SigmaSFR{} relation at any \\Sigmagas.   In this paper, we study the star formation relations and thresholds at 100 pc scale in a sample of low-redshift simulated galaxies. These include simulations representative of Milky Way-like spiral galaxy, the Large and the Small Magellanic Clouds. We analyze the role of interstellar turbulence, gas cooling, and geometry in drawing these relations, by investigating the dependence of the star formation on three parameters: the Mach number, the thickness of the star-forming region and the star formation volume density threshold. We compare the simulated data with the idealized model for star formation of \\cite{Renaud2012}.  Our main findings are as follows: \\begin{enumerate} \\item Our simulations support an interpretation of the surface density threshold for efficient star formation as the typical density for the onset of supersonic turbulence in dense gas, as proposed theoretically by \\cite{Renaud2012}. For all analyzed systems, we obtain qualitatively the same result: regions located below the break are dominated by subsonic turbulence, while turbulence tends to be supersonic in those located in the power-law regime, . \\item The distribution of the ISM of a galaxy in the \\Sigmagas--\\SigmaSFR{} plane (mainly the position of the break) is sensitive to metallicity, but always correlated with the Mach number as detailed above. When different metallicities are considered for otherwise identical systems, \\SigmaSFR{} increases with the metallicity. When different systems with same metallicities are considered (compare Figure \\ref{Fig:contour_all} for \\lmcZsun{} and Figure \\ref{Fig:SMC_metallicity} for \\smcZsun), roughly the same position in the \\Sigmagas--\\SigmaSFR{} plot is obtained. This can explain observations of low-efficiency star formation in relatively dense gas in SMC-like dwarf galaxies. The driving physical parameter is still the onset of supersonic turbulence, but this onset is harder to reach at moderate gas densities in lower-metallicity systems that can preserve warmer gas. \\item The vertical spread in the \\Sigmagas--\\SigmaSFR{} plot is given by the interplay between different parameters of star-forming regions. Figures \\ref{Fig:SMC_model} and \\ref{Fig:MW_model} show a reasonable agreement between simulations and the analytic model of \\cite{Renaud2012}, confirming that this idealized model provides a viable description of star formation in a turbulent ISM compared to more realistic simulations of self-gravitating systems with star formation and feedback. The values of the model parameters (Mach number, thickness and density threshold) characterizing the points in \\Sigmagas--\\SigmaSFR{} plane are close to the values measured in simulations. \\end{enumerate}   Several other models \\citep[e.g.][]{Krumholz2009} have proposed that self-shielding alone is efficient at producing giant molecular clouds and triggering SF. Indeed, this effect cools the gas down at high density, thus enhancing the fragmentation of the ISM, but also lowering the sound speed, i.e. increasing the level of turbulence. Both the compression of the ISM by supersonic turbulence and the fragmenting effect from self-shielding increase with metallicity. Having neglected the dependence of self-shielding on metallicity, our results emphasize only the role of supersonic turbulence in our most metal rich examples. Combining the two effects would lead to a higher efficiency of star formation than either effect alone.  At the scale of clouds, the gravitational collapse is known to trigger SF. However, at larger scales, in galactic structures like spiral arms, we found that the injection of turbulence by self-gravity (and possibly by other processes like shear and feedback) can drive the compression of the gas, leading to SF. In this view, an external trigger like supersonic turbulence could be a sufficient condition to from stars, without necessarily invoking the collapse of large galactic regions ($\\sim$ 100 pc) prior to turbulent compression -- only compressed regions need to eventually collapse into stars."
    },
    "1402/1402.4546_arXiv.txt": {
        "abstract": "We investigate that the two types of the $Q$ balls explain the baryon asymmetry and the dark matter of the universe in the gauge-mediated supersymmetry breaking. The gauge-mediation type $Q$ balls of one flat direction produce baryon asymmetry, while the new type $Q$ balls of another flat direction become the dark matter. We show that the dark matter new type $Q$ balls are free from the neutron star constraint. $n=5$ gauge mediation type and $n=6$ new type $Q$ balls are displayed as an example, where the potential is lifted by the superpotential $\\Phi^n$. These dark matter $Q$ balls may be detected by future observations, such as in advanced IceCube-like observations. ",
        "introduction": "Affleck-Dine baryogenesis is one of the promising mechanism for creating baryon asymmetry of the universe \\cite{AD}. It utilizes a scalar field called the Affleck-Dine (AD) field which carries baryonic charge for producing baryon number when it rotates in the potential. In suspersymmetry (SUSY), there are a lot of flat directions whose scalar potential vanishes in the SUSY limit. They consists of some combination of squarks (and sleptons) \\cite{DRT, GKM}. The scalar potential is lifted by the SUSY breaking effects and the higher order operators coming from the nonrenormalizable superpotential. The field acquires a large field amplitude during inflation, and starts rotation when the Hubble parameter becomes the mass scale of the potential, $H\\sim m_{\\rm eff}$. In the end, the field decays into quarks to be the baryon number of the universe.  Actually, the condensate of the AD field feels spatial instabilities during its rotation, and disintegrated into spherical lumps, $Q$ balls \\cite{KuSh}. The $Q$ ball is a nontopological soliton, the energy minimum configuration of a scalar field whose stability is guaranteed by a finite charge $Q$ \\cite{Coleman}. The charge $Q$ is essentially the baryon number for the AD field. In gauge mediation, large enough $Q$ balls are stable against the decay into nucleons, the lightest particles with unit baryon number, since the mass per charge of the $Q$ ball is less than that of nucleons. In this case, the $Q$ ball can be the dark matter of the universe, while the baryons evaporated from the $Q$ balls would be the baryon number of the universe \\cite {KuSh,KK3}, since all the baryonic charges are absorbed into formed $Q$ balls \\cite{KK1,KK3}.  On the other hand, the $Q$ ball may decay into baryons if the mass per charge of the $Q$ ball is larger than that of the decay baryonic particles. This is achieved when the charge of the $Q$ ball is small enough. In this case, the lightest SUSY particles (LSPs), such as the gravitino \\cite{KK4,KKY, gravitino} or the axino\\cite{KKK}, decayed from the $Q$ ball, may be the dark matter, while the baryons are produced by the $Q$-ball decay as well.  In the former case, however, the direct detection of the dark matter $Q$ balls and the estimation of the cosmological parameters has been refined in this decade since our previous work \\cite{KK3}, we can now conclude that the dark matter $Q$ balls cannot explain the baryon asymmetry of the universe by those baryonic charges evaporated from their surface. See the Appendix for more details.  Then the next question is: Can we still explain the baryon number of the universe by the AD mechanism, while the $Q$ ball is the dark matter? The answer is yes. What we show in this article is that the small enough gauge-mediation type $Q$ balls \\cite{KK3} of one flat direction decay into baryons, and the new type $Q$ balls \\cite{new} of another flat direction are stable to be the dark matter of the universe. The peculiar feature of the dark matter new type $Q$ ball is that it does not suffer from the neutron star constraints, raised in Ref.~\\cite{NS}.  The constructions of the article is as follows. In the next section, we review the features of the gauge-mediation type and the new type $Q$ balls in the gauge mediation. The decay of the gauge-mediation type $Q$-ball is also explained in this section. Abundances of the baryon number and the dark matter in the universe is estimated in Sec.III. We show general constraints on both types of $Q$ balls in Sec.IV, and give an example of $n=5$ unstable gauge-mediation type and $n=6$ stable new type $Q$ balls in Sec.V. We conclude in Sec.VI. In the Appendix, we show the observations exclude the scenario that the dark matter $Q$ ball with evaporated charge being the baryon asymmetry of the universe. ",
        "conclusions": "We have studied the new scenario of the Affleck-Dine $Q$ ball cosmology, in which one flat direction forms unstable $Q$ balls to create baryon number in the universe by their decay, while the other direction builds up into stable $Q$ balls to be the dark matter of the universe. Since the gravitino mass is usually smaller than that of nucleons in the gauge-mediated SUSY breaking, the new type $Q$ ball is always stable against the decay into nucleons. At the same time, the gauge-mediation type $Q$ ball with small enough charge $Q$ can decay into nucleons, releasing all the baryon number they had contained.  We have shown that this scenario works well for wide range of the parameter space: $M_F=10^5 - 10^7$~GeV and $Q_{\\rm G} = 10^{17} - 10^{25}$ for the gauge-mediation type with $m_{3/2} \\lesssim$~GeV and $Q_{\\rm N} \\gtrsim 10^{25}$ for the new type. We have seen that this dark matter new type $Q$ balls could be detectable in the (future) IceCube-like experiments.  We have also presented more concrete model: the $n=5$ direction forms gauge-mediation type $Q$ balls, while the $n=6$ direction becomes the new type $Q$ balls. In this case, we can pin down the model parameters more accurately.  In addition, the astrophysical constraints from the neutron star \\cite{NS} do not exist. Since the $Q$ balls form from the rotating AD condensate, and its rotation is due to the $A$ terms in the potential, the same $A$ terms prevent the $Q$ balls captured by the neutron star from growing up at the certain charge. Thus, they would not swallow the baryonic charge of the neutron star so much, and are even destroyed by the effects of the $A$ terms. Therefore, the new type $Q$ ball can safely be the dark matter of the universe."
    },
    "1402/1402.6891_arXiv.txt": {
        "abstract": "We present a detailed study of the interaction process of two coronal mass ejections (CMEs) successively launched on 2011 February 14 (CME1) and 2011 February 15 (CME2). Reconstructing the 3D shape and evolution of the flux ropes we verify that the two CMEs interact. The frontal structure of both CMEs measured along different position angles (PA) over the entire latitudinal extent, reveals differences in the kinematics for the interacting flanks and the apexes. The interaction process is strongly PA-dependent in terms of timing as well as kinematical evolution. The central interaction occurs along PA-100$^{\\circ}$, which shows the strongest changes in kinematics. During interaction, CME1 accelerates from $\\sim$400~km~s$^{-1}$ to $\\sim$700~km~s$^{-1}$ and CME2 decelerates from $\\sim$1300~km~s$^{-1}$ to $\\sim$600~km~s$^{-1}$. Our results indicate that a simplified scenario like inelastic collision may not be sufficient to describe the CME-CME interaction. Magnetic field structures of the intertwining flux ropes as well as momentum transfer due to shocks play an important role in the interaction process. ",
        "introduction": "The kinematical behavior of isolated coronal mass ejections (CMEs) can strongly differ compared to a series of CMEs. Successive eruptions from the Sun cause a change of the interplanetary environment (e.g. density, speed) and complex magnetic structures may form during CME-CME interaction. Observable features of CME-CME interaction are restricted to white-light and radio data. The CME kinematics is found to clearly change during the collision of two CMEs, resulting in either strong deceleration \\citep[e.g.,][]{temmer12} or acceleration \\citep[][super-elastic collision]{shen12}. Strong deceleration in the early kinematical evolution was also reported for CMEs that run into strong overlying magnetic fields \\citep{temmer10}. It is generally accepted that the shape and orientation of the magnetic flux rope embedded in a CME plays an important role for the further propagation and evolution in interplanetary space \\citep[e.g.][]{schmidt04, lugaz05, lugaz13}.   The kinematical behavior for a chain of CME events from 2011 February 14 to 2011 February 15 and their production of a Forbush decrease is described by \\cite{maricic13}. They find evidence for a gradual momentum transfer between the slower and the faster CME, resulting in the deceleration of the faster and the acceleration of the slower one \\citep[see also][]{farrugia04}. The momentum transfer may take place through a CME-driven shock. This interpretation is supported by numerical simulations showing that fast shocks may travel through slower CME fronts ahead \\citep[see e.g.,][]{lugaz09,liu12}.    The current analysis is an extension of the study by \\cite{maricic13}. We investigate in detail the interaction process of the two consecutive CMEs launched in a time window of about 8 hours during 2011 February 14--15. Using spherically de-projected Heliospheric Imager data in combination with a 3D reconstruction of the magnetic flux ropes, we track the frontal structures of both CMEs over their entire latitudinal extent in order to study asymmetries in the interaction process. We also investigate dynamic radio spectra in the decametric/hectometric domain to study the nature of the interaction process. ",
        "conclusions": "We followed in detail the CME-CME interaction event from 2011 February 14--15. Applying a graduated cylindrical shell model \\citep{thernisien06} to the multi-spacecraft coronagraphic data we derive that both CMEs head into the same direction, with different speed, size and axial orientation. The front of CME1 is less bright than CME2 and its structure is intermittent and increased in intensity close to PA-70$^{\\circ}$ and PA-100$^{\\circ}$, which matches with the location where both flux ropes are supposed to interact.   From the kinematical curves measured along different PAs (cf.\\ Figure~\\ref{resid}), we derive different speeds and different times for the encounter between the fronts of both CMEs. This might be a consequence of the different axial orientations of the flux ropes estimated to deviate from each other by $\\sim$35$^{\\circ}$ and would indicate that the interacting magnetic field structures are decisive for the evolution of the interaction process. We would like to note that the strong kinematical changes measured for the frontal structures along PA-80--90$^{\\circ}$, affect both CMEs at about the same time. Since CME2 propagates in the wake of CME1, changes in the speed of CME1 most likely affect the kinematics of CME2. Inspecting the behavior of the southern part of CME2 along PA-125$^{\\circ}$, we obtain that its frontal structure is highly deformed and propagates with much higher speed (960~km~s$^{-1}$) than the frontal structure measured along PA-70--100$^{\\circ}$ (530$\\pm$50~km~s$^{-1}$). These results give further evidence on the latitudinal range of influence of the flux rope of CME1, associated with strong changes of the ambient density, speed and magnetic flux. Deformation of CME fronts (concave-outward shape, pancaking) due to different solar wind speeds in low and high latitudinal regions is reported by e.g., \\cite{russell02}, \\cite{riley04}, and \\cite{liu06}.   For the direction of central collision along PA-100$^{\\circ}$, we perform the exercise of assuming the simplified scenario of an inelastic collision between two bodies of similar mass (cf.\\ Table~\\ref{table1}). To push CME1 from 400 to 700~km~s$^{-1}$, as observed around $\\sim$9~UT, CME2 would need a speed of at least 1000~km~s$^{-1}$. CME2 reached a maximum speed of $\\sim$1300~km~s$^{-1}$ within the FoV of COR2, but decelerated during the assumed time of collision already to $\\sim$600~km~s$^{-1}$. This indicates that CME2 may influence CME1 much earlier than suggested by the white-light data. Similar results are found for the CME-CME interaction from 2010 August 1, for which a strong deceleration of the faster CME was observed a few hours before its leading edge showed in the heliospheric images the merging with the slower CME ahead \\citep{temmer11}. \\cite{maricic13} concluded that the CMEs are of finite thickness, thus to a finite ``signal'' travel time the interaction process may start well before changes are observed at the leading edge of CME1, most likely through transfer of momentum \\citep[see e.g.,][]{farrugia04}. A recent paper by \\cite{lugaz13} identifies a number of problems in determining the type of collision between two CMEs among others due to changes in the mass of the CMEs during evolution \\citep[see also][]{bein13} as well as ongoing perturbation hours after the collision.   Several enhancements of the radio type II burst associated to CME2 are observed and may give hint that the interaction process started as early as 2011 February 15 02:40~UT \\citep[see e.g.,][]{gopalswamy01}. However, this first enhancement appears about 1 hour after the launch of CME2, much too early for CME2 or the shock of CME2 ($v_{\\rm shock}\\approx$ 1100~km~s$^{-1}$) to reach CME1 which is roughly at a distance of 20~R$_{\\odot}$ at that time. Recent findings show that radio enhancements in type II bursts are frequently related to shock-streamer interactions \\citep{shen13}. As a large part of CME2 (PA-70--100$^{\\circ}$) propagates in the rear of CME1, the radio enhancements could be a consequence of the interaction between the shock of CME2 and the streamer-like post-eruption current sheet formed behind CME1. Interestingly, the enhancements are found to be associated to type III bursts which are all related to flares emitted from AR~11158, the same AR where both CMEs are launched from. Type III radio bursts subsequently observed on 2011 February 15 (not shown), stemming from the same AR, stop at frequencies related to the downstream region of the extrapolated type II burst, as if it would be a barrier for particles entering the magnetic structure \\citep[see also][]{macdowall89}. In a follow-up study we aim to make a more thorough analysis on the connection between the cutoff frequencies of type III and type II bursts as well as their relation to the process of CME-CME interaction. In conclusion, our observational data reveal only the consequences of CME-CME interaction that may not be sufficient to fully describe the process of interaction. The shape and orientation of the magnetic structures of CME-CME events certainly play a key role."
    },
    "1402/1402.2360_arXiv.txt": {
        "abstract": "Full polarization measurements of 1665 and 1667-MHz OH masers at sites of massive star formation have been made with the Parkes 64-m radio telescope.  Here we present the resulting spectra for 104 northerly sources, from Galactic longitude 350$^\\circ$ through the Galactic Centre to 41$^\\circ$. Some maser positions were previously uncertain by many arcseconds, and thus for more than 20 masers we made new measurements with the ATCA (which also revealed several hitherto unreported masers), in most cases yielding arcsecond  precision to match the majority of sites. Position improvements have assisted us  in distinguishing OH masers with accompanying methanol masers from those without (thought to be at a later stage of evolution). There was no existing linear polarization information at many sites, and spectral resolution was sometimes poor, or velocity coverage incomplete. These inadequacies are addressed by the present Parkes spectra. The whole OH maser sample exhibit the well-known predominance of highly circularly polarized features.  We find that linear polarization is also common, but usually much weaker, and we highlight the rare cases of very pronounced linear polarization that can extend to 100 per cent. Unusually large velocity ranges of at least 25 \\kms\\ are present at seven sites. Our spectra measurements for most sources are at two epochs spaced by nearly one year, and reveal high stability at most sites, and marked variability (more than factors of two in the strongest feature) at only five sites.  The spectra also provide a valuable reference for longer term variability, with  high stability evident over the past decades at 10 sites and  marked variability for four of the sample. Future systematic monitoring of these variables may uncover further examples of periodicity, a phenomenon so far recognised in only one source. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.7129_arXiv.txt": {
        "abstract": "We report on PTF1\\,J191905.19+481506.2, a newly discovered, partially eclipsing, outbursting AM CVn system found in the Palomar Transient Factory synoptic survey. This is only the second known eclipsing AM CVn system. We use high-speed photometric observations and phase-resolved spectroscopy to establish an orbital period of $22.4559(3)\\,$min. We also present a long-term light curve and report on the normal and super-outbursts regularly seen in this system, including a super-outburst recurrence time of $36.8(4)\\,$d. We use the presence of the  eclipse to place upper and lower limits on the inclination of the system and discuss the number of known eclipsing AM CVn systems versus what would be expected. ",
        "introduction": "AM CVn systems are rare, ultra-compact, semi-detached, white dwarf binaries with periods ranging from 5 to 65 minutes. First identified over 40 years ago by \\citet{1967AcA....17..255S}, it has only been in the last decade that the number of known systems has risen above ten. Yet despite the recent discovery of almost 25 additional systems, their rich and complex phenomenological behavior has limited our understanding of this class of post-period minimum binaries. They are considered to be the helium analog of cataclysmic variables (CVs) and are important as strong, low-frequency Galactic gravitational wave sources \\citep{2004MNRAS.349..181N,2007MNRAS.382..685R,Nissanke:2012fk}, the source of the proposed ``.Ia'' supernovae \\citep{2007ApJ...662L..95B}, and one of the believed end-points of binary white dwarf evolution \\citep{2001A&A...368..939N}. However, many questions about these unique systems remain, including their population density, their evolutionary pathways, and the interactions between the two components, including the He-rich accretion disk. We refer the reader to \\citet{2005ASPC..330...27N} and \\citet{2010PASP..122.1133S} for general reviews.  Follow-up study of a newly-identify AM CVn system typically begins with a spectroscopic determination of its orbital period using phase-resolved spectroscopy of the ``hot spot'' \\citep{1981ApJ...244..269N}, which is the location where matter from the donor hits the accretion disk. However, these measurements are incomplete, as the inclination of the systems cannot be determined. Only those systems showing eclipses and radial velocity variations of two components (e.g., the donor and accretor or the accretor and hot spot) allow for a full determination of the system's parameters (e.g., component masses, inclination, period, etc). Eclipses also allow for the extremely precise measurement of the orbital period change. For example, the orbital period of the double detached white-dwarf binary SDSSJ0651+2844 was measured to be decreasing by \\citet{Hermes:2012ys} while the only known eclipsing AM CVn system, SDSSJ0926+3624 (\\citealt{2005AJ....130.2230A}; \\citealt{2011MNRAS.410.1113C}, hereafter C11), was recently shown to have an increasing orbital period \\citep{2013arXiv1309.4087S}.  AM CVn systems have been observed to have three relatively distinct phenomenological states. Systems with orbital periods (hereafter $P_{orb}$) less than $20\\,$min have been observed in a ``high'' state, characterized by optically thick accretion disks and absorption-line spectra. In contrast, systems with $P_{orb}>40\\,$min, are observed to have emission-line spectra from what are believed to be optically-thin accretion disks and are said to be in quiescence (though see \\citealt{Woudt:2013kx} for an example of a longer period system that has been observed to outburst). Systems in neither class show photometric variability over 0.5\\,mag.  Between these orbital period limits are outbursting systems, which are observed to have dwarf nova-type outbursts of 3--6\\,mag as well as variability at the 10\\% level in both quiescence and outburst \\citep[e.g., ][]{1997PASP..109.1100P}. Recent studies have shown that the frequency of these outbursts decreases as the orbital period increases (\\citealt{2011ApJ...739...68L}; \\citealt{2012MNRAS.419.2836R}).  The population density of AM CVn systems has not been conclusively determined: population synthesis estimates of the space density \\citep{2001A&A...368..939N} have not been observationally confirmed by color-selected samples \\citep{2007MNRAS.382..685R,Carter:2013fk}. The latest results of \\citet{Carter:2013fk} suggest a space density of $(5\\pm3)\\times10^{-7}\\,\\text{pc}^{-3}$, a factor of 50 lower than the population synthesis estimates by \\citet{2001A&A...368..939N}. The reason for this discrepancy is currently unknown. We note that these color-selected samples are mostly sensitive to longer period systems.  Over the last two years, we have conducted a search for outbursting AM CVn systems using the Palomar Transient Factory\\footnote{\\url{http://ptf.caltech.edu/}} \\citep[PTF; ][]{2009PASP..121.1395L,2009PASP..121.1334R} large-area synoptic survey, in part, to use a different approach to the discovery of AM CVn systems that does not rely on their colors. The PTF uses the Palomar $48^{\\prime\\prime}$ Samuel Oschin Schmidt telescope to image up to $\\sim2,000\\deg^{2}$ of the sky per night to a depth of $R\\sim20.6$ or $g'\\sim21.3$. After identifying outbursting systems, we obtain classification spectra. To date, this survey has identified $\\mathord{>}340$ cataclysmic variables (CVs), six new AM CVn systems, and one extremely faint AM CVn candidate \\citep{2011ApJ...739...68L,Levitan:2013uq}.   In this paper, we present PTF1\\,J191905.19+481506.2 --- a new AM CVn system discovered by the Palomar Transient Factory. This system is particularly interesting because: \\begin{itemize} \\item shallow eclipses are present, making it only the second known eclipsing AM CVn system. \\item its orbital period is the shortest known of outbursting AM CVn systems. \\end{itemize} We note that while this system is in the Kepler field, it, unfortunately, fell into a gap between the detectors of the \\textit{Kepler} satellite.  This paper is organized as follows. In Section \\ref{sec:detection} we discuss our data reduction and analysis methods. In Section \\ref{sec:observations} we present both short-term and long-term photometric and spectroscopic data and perform a period analysis. We consider the system's geometric structure and the rate of eclipsing AM CVn systems in Section \\ref{sec:discussion}. We summarize in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have established PTF1J1919+4815 as an eclipsing AM CVn system from its He-rich and H-deficient spectrum, its phenomenological behavior, and its photometrically determined period. Its orbital period was measured to be $22.4559(3)\\,$min, but we did not find a significant hot spot in the S-wave and Doppler map. PTF1J1919+4815 shows a well-defined super-outburst that is $36.8\\pm0.4\\,$d. Such a non-variable recurrence time is atypical for AM CVn systems. We used these measurements to constrain the inclination of the systems given assumptions about its evolution and to consider the geometric structure of the system."
    },
    "1402/1402.0015_arXiv.txt": {
        "abstract": "The spectroscopic studies of near infrared emission arising from supernovae allow to derive crucial quantities that could better characterise physical conditions of the expanding gas, such as the CaII IR HVF spectral feature. For this reason is mandatory to have Diffractive Optical Elements (DOEs) with a spectral coverage in the range 8000 - 10000 $\\textrm{\\AA}$ (for low z sources) combined with a reasonable Signal to Noise Ratio (S/N) and medium-low resolution. In order to cope with all of those requirements we developed a Volume Phase Holographic Grating (VPHG) based on an innovative photosensitive material, developed by Bayer MaterialScience. We demonstrated the capabilities of this new DOE through observation of SN 2013fj as case study at Asiago Copernico Telescope where AFOSC spectrograph is available. ",
        "introduction": "Astronomical spectrographs are key instrumentation to tackle the open issues in astronomy. One of the most important element along with the detector is the dispersing element. In the last 15 years, the Volume Phase Holographic Grating (VPHG) technology has gained a lot of interest in astronomical field and it has been used in some spectrographs \\citep{baldry, andrea10, barden, pazder}. The reasons for such interest stem from the fact that these gratings show unique features, such as i) the high peak efficiency (up to 100\\% theoretically) both at low and large dispersion; ii) the ease of performance tuning and customisation (each VPHG is a master grating). A VPHG consists in a thin layer of holographic material, usually dichromated gelatine \\citep{andrea10, barden}, which is sandwiched between two glass windows. The phase of incident light is modified passing through the gratings thanks to a periodic modulation (usually sinusoidal) of the refractive index written in the holographic material. Hence the fundamental parameters that rule the overall efficiency of such devices are the thickness of the active film and the modulation of the refractive index inside it. VPHG technology has been applied in different astronomical spectrographs, at room and cryogenic temperatures \\citep{wiyn, andrea7, andrea3, lepine, hou, andrea1, andrea9}. They have also been used as tuneable filters and cross dispersers \\citep{mendes, gibson, castilho}. The VPHGs assembled in a GRISM configuration have been also used in this kind of instrumentation. The availability of a low resolution grism covering the red part of the optical spectrum is of paramount importance for several astrophysical targets, among which stands out the study of supernovae in general and type Ia supernovae in particular. Such a grism will characterise the evolution of important lines during the photospheric (e.g. O I 7774 $\\textrm{\\AA}$ and Ca II IR triplet), and the nebular phases (e.g. Ca II] 7291-7323 $\\textrm{\\AA}$ and Ca II IR triplet).\\\\ Particularly important will be to study the evolution in type Ia supernovae of the high velocity features (HVFs) seen in the Ca II triplet profile, especially at early, pre-maximum, phases \\citep[see][for a recent review]{chil13}. The origin of the HVFs remains unknown, but different hypotheses have been proposed, which include the impact of the ejecta with a circumstellar shell lost by the progenitor before the explosion \\citep[see][]{ger2004}; an enhancement in the abundance of intermediate-mass elements (IMEs) in the outermost layers of SN Ia ejecta \\citep{mazzali05a, mazzali05b, tanaka08}; or variations in the ionisation state of IMEs in the outer layers of SN Ia ejecta \\citep{blondin2013}.\\\\ It is evident that the study of the Ca II HVFs could have a big impact in deriving the true progenitor scenario involved in the SNIa explosion, which in turn could have important implication in the use of SNIa in Cosmology. Moreover, the circumstellar (CS) material responsible for the HVFs could cause subtle alteration of the spectral energy distribution of SNIa, which could have possible consequences in the luminosity standardization of SNIa \\citep{chil13}. \\\\ In order to accomplish the desired requirements we designed and manufactured a VPHG based on a completely new holographic material. The study of this new material, other than dichromate gelatins (DCGs), is very important in order to make possible the design and manufacturing of innovative and large VPHGs. Indeed DCGs are difficult to handle, they require a complex chemical process and the scalability to very large size gratings can be an issue. For this reasons we focused the attention to solid photopolymers  which combines high throughput, high refractive index modulation, self developing (i.e. no chemical processes needed) and size scalability. Such new material belongs to the class of solid photopolymers and this is the first time, in our knowledge, that this kind of holographic material has been used to make scientific grade dispersing elements. In the past only liquid photopolymers have been used once to produce VPHGs mounted in MOIRCS and FOCAS \\citep{andrea4, andrea6}. Starting from the scientific case of the SN 2013fj hereinafter we demonstrate the capabilities in terms of spectral resolution and throughput of this new family of Volume Phase Holographic Grating, based on  photopolymers, comparing two spectra of SN 2013fj in which one of them is secured by the adoption of previous existing state of the art grating. Throughout this paper we refer to this new grating as VPH6. ",
        "conclusions": "\\label{sec:discussion}  We demonstrated the good performances obtainable by using a Volume Phase Holographic Grating based on new photopolymer materials which are self developing, characterised by a high sensitivity and dynamic range in conjunction with an easy processability. The GRISM has been designed and manufactured in order to maximise the efficiency reducing the reflection losses. For these reasons such devices are a reliable alternative to classical VPHGs. We assessed the scientific requirements which drawn the design of this new DOE by collecting the spectrum of the newly discovered SN Ia PSN J22152851+1534041 = SN 2013fj . We finally carried out a sound comparison with another spectrum of the same object under the same conditions, secured with the standard grism that is characterised by the same dispersion and resolution. \\\\ \\\\ \\textbf{Acknowledgements} \\\\ We are grateful to: Dr. Thomas F\\\"{a}cke of Bayer MaterialScience for providing the material and for the useful discussions; The technicians at mt. Ekar for all the support during commissioning based on observations collected at Copernico telescope (Asiago, Italy) of the INAF - Osservatorio Astronomico di Padova; L.T. and S.B. are partially supported by the PRIN-INAF 2011 with the project \"Transient Universe: from ESO Large to PESSTO\". \\newpage \\clearpage"
    },
    "1402/1402.0824_arXiv.txt": {
        "abstract": "{The impact of new stellar evolution models with rotation on the predictions of population synthesis models is discussed. Massive rotating stars have larger convective cores than their non-rotating counterparts,  and their outer layers are chemically enriched due to increased mixing. Together, these two effects lead to hotter and more luminous stars, in particular during later evolutionary phases. As a result, stellar populations containing massive stars are predicted to become more luminous for a given mass and to emit more ionizing photons. Depending on the assumed rotation velocity, rotation causes profound changes in the properties of young stellar populations. These changes are most noticeable at later evolutionary phases and at shorter wavelengths of the spectral energy distribution. Most strikingly, the Lyman continuum luminosity increases by up to a factor of five in O- and Wolf-Rayet stars. Care is required when comparing these models to observations, and some fine-tuning of the models is still required before recalibrations of star-formation indicators should be attempted. } ",
        "introduction": "Population synthesis models aim to predict photometric, spectroscopic, and chemical properties of stellar systems, such as massive star clusters or entire galaxies. Stellar evolution models are at the heart of population synthesis, as these models provide the prescription for the relation between stellar luminosity and mass and their evolution with time. Due to the time scales involved, stellar evolution mostly eludes direct observation and is therefore heavily model-dependent. Nevertheless, great strides have been made over the past decades in the field of massive-star evolution (e.g., \\cite{ChiMae86}; \\cite{MaeMey00}; \\cite{Lan12}). Each new generation of evolution models has led to transformational changes in the field of population synthesis. The present epoch marks another such transformation: stellar rotation, which had previously been recognized to be significant for the evolution of massive stars, has finally been implemented self-consistently in evolutionary tracks. In this contribution I will introduce a new set of evolutionary tracks for massive stars that accounts for rotation and discuss how these tracks affect population synthesis models. ",
        "conclusions": ""
    },
    "1402/1402.1633_arXiv.txt": {
        "abstract": "We report the Suzaku/XIS results of the Galactic oxygen-rich supernova remnant (SNR), G292.0$+$1.8, a remnant of a core-collapse supernova. The X-ray spectrum of G292.0$+$1.8 consists of two type plasmas, one is in collisional ionization equilibrium (CIE) and the other is in non-equilibrium ionization (NEI). The CIE plasma has nearly solar abundances, and hence would be originated from the circumstellar and interstellar mediums. The NEI plasma has super-solar abundances, and the abundance pattern indicates that the plasma originates from the supernova ejecta with a main sequence of 30--35\\,{\\it{M}}$_{\\odot}$.  Iron K-shell line at energy of 6.6\\,keV is detected for the first time in the NEI plasma. ",
        "introduction": "X-ray spectra of optically-thin hot plasmas in supernova remnants (SNRs) provide key information on the nucleosynthesis during the stellar  evolution and the supernova (SN) explosion. Since iron (Fe) is the final product of the major nuclear reaction network in a massive star and a SN, it is particularly important element.  So far, Fe abundances of some SNRs have been estimated using the Fe L-shell lines, which are dominant in the low energy band around 1\\,keV.  This energy band is complex with other strong emission lines such as oxygen (O) and neon (Ne). Moreover the Fe L-lines consist of many emission lines in this narrow energy band, and hence conventional X-ray detectors such as X-ray CCD cannot resolve these many lines. The Fe K-shell lines at 6.4--6.7\\,keV are simpler and their emission model is more reliable than those of the L-shell lines. However, due to the limited line fluxes and the sensitivity in the energy band above 6\\,keV, the Fe K-shell lines have been reported from only a limited number of X-ray bright SNRs.  Core-collapse (CC) SNe eject less Fe than those of Type Ia SNe (\\citealt{Nomoto1984}; \\citealt{Iwamoto1999}, and references therein).  Most of the Fe in the center of the progenitor star is collapsed into a central compact object (neutron star or black hole). Thus observationally,  Fe K-shell lines would become weaker than those in Type Ia. The flux of Fe may come from the combined effect of the elemental abundances and the thermal state of the gas.  Thus the differences between Type Ia and CC SNRs may be related to the initial condition of the circumstellar medium (CSM) and an explosion mechanism.  Fe K-shell lines have been reported from other candidates of CC SNRs (e.g., W49B: \\citealt{Ozawa2009}; IC443: \\citealt{Yamaguchi2009}; G349.7+0.2: \\citealt{Slane2002} and G350.1-0.3: \\citealt{Gaensler2008}; 3C397: \\citealt{Chen1999}). However, some of them are controversial whether they are CC SNRs or not.  The most reliable criteria of CC SN are the presence of a neutron star (pulsar), a pulsar wind nebula (PWN), and oxygen-rich (O-rich) knots, because O is largely enhanced in the ejecta of a CC SN.  At present, three O-rich remnants, Cas A, Puppis\\,A and G292.0$+$1.8, have been reported in our Galaxy (see \\citealt{Vink2012}, and references therein).  Cas A is the most luminous, and hence the Fe K-shell line was firstly detected at 6.6\\,keV (\\citealt{Tsunemi1986}). The spatial distribution is not a simple stratified structure (\\citealt{Willingale2002}). Thus the SN  explosion would be highly asymmetric. Puppis\\,A has also the interesting distribution of the ejecta, in that the ejecta is found only in the east, mostly north-east portion (\\citealt{Winkler1985}; \\citealt{Hwang2008}; \\citealt{Katsuda2010}), while a neutron star is propelled in the opposite direction (e.g., \\citealt{Becker2012} and references therein). Such a recoil between SN ejecta and a neutron star is expected in a recent SN explosion model (\\citealt{Scheck2006}), and would be examined by asymmetric Fe distribution.  However, no Fe K-shell line has been observed from Puppis\\,A.  For the study of the Fe K-shell line in CC SNRs, we observed G292.0$+$1.8. The discovery of a pulsar and a PWN in G292.0$+$1.8 is further confirmation of a CC SNR (\\citealt{Hughes2001}, \\citeyear{Hughes2003}; \\citealt{Camilo2002}; \\citealt{Gaensler2003}). The pulsar is located about 0.9\\arcmin southeast from the geometrical center of the SNR.  The distance of G292.0$+$1.8 is estimated to be  6\\,kpc  (\\citealt{Gaensler2003}),  and  the  age  is likely 2990$\\pm$60\\,years (\\citealt{Winkler2009}).  The morphology of G292.0$+$1.8 consists of many small knots and the central belt-like filaments running from the east to the west (\\citealt{Park2002}). The central filaments have a normal solar-type composition, suggesting that these are the shocked CSM.  \\citet{Lee2010} reproduced the intensity profile of the outer CSM region by a slow wind from a red supergiant (RSG) star with the total mass of the wind of 15--40\\,{\\it{M}}$_\\odot$. The implied progenitor mass ($M >$ 20\\,{\\it{M}}$_\\odot$) was in plausible agreement with previous estimates (\\citealt{Hughes1994}; \\citealt{Gonzalez2003}; \\citealt{Park2004}). The knots have an enhanced metallicity; Si is enhanced in north-northeast, O is enhanced primarily in southeast, Ne is in northwest and southeast, and Mg is in northwest (\\citealt{Park2002}). These knots are probably ejecta origin. The asymmetric distribution of the ejecta elements is interpreted to be non-uniform thermodynamic conditions of the X-ray-emitting ejecta (\\citealt{Park2007}).  In spite of these extensive studies, no significant Fe K-shell line has been detected.  \\citet{Park2004} proposed that the ejecta are strongly stratified by composition and the reverse shock has not propagated to the Fe rich-zone yet.  However, the X-ray spectra reported so far are faint in the hard band (except for that of PWN), and hence observed lines have been limited up to sulfur (S) K-shell line.  In this paper, we report the Suzaku discovery of an Fe K-shell line in a high temperature plasma ($kT_e$ = 2--3\\,keV) extending to $E$ = 10\\,keV.  K-shell lines of argon (Ar) and calcium (Ca) are also reported.  Based on the wide band spectral analysis, we discuss the nature of G292.0$+$1.8. We adopt the solar abundances of \\citet{Anders1989}. Unless otherwise specified, all errors represent 1$\\sigma$ confidence levels. ",
        "conclusions": "In the plasma evolution in SNRs, the X-ray emissions have two different components: the CSM (plus ISM) heated by the blast wave, and the ejecta from the progenitor star heated by the reverse shock.  In the spectral fitting of G292.0$+$1.8, we find two type plasmas in CIE and NEI conditions; the 2-APEC ($kT_e\\sim$\\,0.2 and 0.7\\,keV) and the 2-VPSHOCK ($kT_e\\sim$\\,1 and 2.5\\,keV) plasmas. We call these two type plasmas, the low-$kT_e$ plasma and the high-$kT_e$ plasma, respectively. Since the low-$kT_e$ has nearly solar abundances for all elements and the high-$kT_e$ has super-solar abundances (see table \\ref{SNR-PWN_para}), these would be the CSM plus ISM and the ejecta origin of a CC SN, respectively. Chandra spectra from many selected regions of bright small spots are described by 1-VPSHOCK model with super-solar abundances (\\citealt{Park2004}), while those from the faint outer-most shell are 1 or 2-VPSHOCK model with sub-solar to solar abundances (\\citealt{Gonzalez2003}; \\citealt{Lee2010}). These results of \"no-CIE\" plasma are in contrast to the existence of CIE ($\\sim$ solar) components in the Suzaku spectra. Since the Chandra results are from selected spots or filaments and those of Suzaku are from the entire SNR region, we suspect that our CIE plasmas are prevailing over the entire SNR except the outer-most shell, while the bright spots are dominated by the VPSHOCK plasma.  We discover Fe K-shell line at 6.6 keV in the eject plasma for the first time. The energy indicates that ionization state of Fe is around B-like.  This medium ionization state is similar to another young CC SNR, Cas A, but is in contrast to nearly Ne-like states in young well known Type Ia SNRs, Tycho, Kepler, and SN\\,1006.  Figure \\ref{ejecta_abundance_pattern} is the abundance plot of the ejecta for O, Ne, Mg, Si, S, Ar and Fe relative to Si (from table \\ref{SNR-PWN_para}) together with  those of the CC SN model in various progenitor masses (\\citealt{Woosley1995}). We see that the observed abundance pattern is globally in agreement with the model of 30--35\\,{\\it{M}}$_\\odot$.  These mass range of the progenitor star confirm the previous report of 30--40\\,{\\it{M}}$_\\odot$ (\\citealt{Gonzalez2003}), which was estimated based on the limited information of non-detection of the explosively synthesized heavy elements such as Ar, Ca and  Fe. One may argue that CC SN of a massive progenitor 30--35\\,{\\it{M}}$_\\odot$ may leave a black hole rather than a neutron star.  However, other observations show that a neutron star can be still formed from even these massive progenitor stars (\\citealt{Muno2006}).  Although \\citet{Park2004} and \\citet{Gonzalez2003} reported  significant spatial variations in the sub-arcmin scale, we find no significant and systematic variations in the spatial scale over arcmin.  In fact, the best-fit abundances of most of the heavy elements show no variations within their typical errors of 50\\%. Nevertheless, we find marginal evidence of spatial variation of Fe in the ejecta; the north region is enhanced compared to that of the center region. Since the position of the neutron star (PWN) is off-set to southeast from the geometrical SNR center (\\citealt{Park2007}), it would be conceivable that Fe from the core region would be ejected  to the opposite northwest direction. Our observational result of the Fe variation is marginal to support this off-set effect due to large errors. To establish this kick-off scenario, we need higher quality observations.  The best-fit spectral parameters of the PWN, the photon index and unabsorbed flux (4--8\\,keV) are 1.91$\\pm$0.03 and (3.80$\\pm$0.18)$\\times$10$^{-4}$ photons s$^{-1}$ cm$^{-2}$, respectively. The photon index is steeper than that of the pulsar (1.6--1.7, \\citealt{Hughes2001}, \\citeyear{Hughes2003}). Probably the index increases as the distance from the central pulsar increases (e.g. \\citealt{Slane2000}). The PWN flux is 52\\% of the total flux (4--8\\,keV) from the whole SNR ((7.30$\\pm$0.06)$\\times$10$^{-4}$ photons s$^{-1}$ cm$^{-2}$). This ratio is slightly smaller than 66\\%, determined with the high spatial resolution observation of Chandra (Hughes et al. 2001). This difference, however,  would be within uncertainty range due possibly to the NXB and CXB subtraction\\footnote{The 4-8 keV band fluxes of the BG region are about 8\\% and 70\\%  of the whole SNR regions for the data of Suzaku and Chandra, respectively}, and/or other systematic cross errors including different data reduction processes between Suzaku and Chandra. Thus our simultaneous fitting analysis properly estimates the flux and spectra of both the SNR and the PWN, although the spatial resolution of Suzaku is limited to completely separate the emissions from these two sources."
    },
    "1402/1402.3293_arXiv.txt": {
        "abstract": "We explicitly test the equal-time consistency relation between the angular-averaged bispectrum and the power spectrum of the matter density field, employing a large suite of cosmological $N$-body simulations. This is the lowest-order version of the relations between $(\\ell+n)$-point and $n$-point polyspectra, where one averages over the angles of $\\ell$ soft modes. This relation depends on two wave numbers, $k'$ in the soft domain and $k$ in the hard domain. We show that it holds up to a good accuracy, when $k'/k\\ll 1$ and $k'$ is in the linear regime, while the hard mode $k$ goes from linear ($0.1\\,\\hmpci$) to nonlinear ($1.0\\,\\hmpci$) scales. On scales $k\\lesssim 0.4\\,\\hmpci$, we confirm the relation within the statistical error of the simulations (typically a few percent depending on the wave number), even though the bispectrum can already deviate from leading-order perturbation theory by more than $30\\%$. We further examine the relation on smaller scales with higher resolution simulations. We find that the relation holds within the statistical error of the simulations at $z=1$, whereas we find deviations as large as $\\sim 7\\%$ at $k \\sim 1.0\\,\\hmpci$ at $z=0.35$. We show that this can be explained partly by the breakdown of the approximation $\\Omegam/f^2\\simeq1$ with supplemental simulations done in the Einstein-de Sitter background cosmology. We also estimate the impact of this approximation on the power spectrum and bispectrum. ",
        "introduction": "The large-scale structure of the Universe provides us with a wealth of information on the initial conditions of the Universe as well as the underlying gravity theory that governs the time evolution on sufficiently large scales \\cite{Peebles1993,Laureijs2011}. A classic tool for discussing its statistical properties are the polyspectra of the matter density field at a given time (the Fourier transforms of the $n$-point correlation functions) \\cite{Peebles1980,Bernardeau02}. The power spectrum, the lowest-order polyspectrum, has played a central role to test cosmological models and determine their parameters precisely. Standard models of the early Universe predict almost Gaussian initial conditions, in agreement with a number of observational probes (e.g., measures of cosmic microwave background anisotropies \\cite{Planck_nG}). However, even if the initial conditions are perfectly Gaussian, the cosmic density field at late times exhibits non-Gaussian features acquired through the nonlinear gravitational dynamics.  The polyspectra induced by gravity can be analytically derived order by order using standard perturbation-theory techniques (see Ref.~\\cite{Bernardeau02} for a review). In these calculations, an approximate treatment is usually adopted that greatly simplifies the structure of the basic equations. That is, the combination $\\Omegam/f^2$ is replaced with unity, where $\\Omegam$ is the time-dependent matter density parameter and $f\\equiv\\mathrm{d}\\ln D_+/\\mathrm{d}\\ln a$ is the linear growth rate, with $D_+$ being the linear growing mode. This approximation is exact in the Einstein-de Sitter universe and sufficiently accurate in most other cosmological models based on general relativity, because i) one usually recovers Einstein-de Sitter at early times and ii) over the realistic range of cosmological parameters one has $f \\simeq \\Omega_{\\rm m}^{\\gamma}$ with $\\gamma \\simeq 0.5$ \\cite{Peebles1980}. When this approximation is applied, all the dependence on the cosmological parameters is absorbed by the linear growth rate $D_+$, and the time dependence of the solution is also fully encapsulated in $D_+$. This simplifies perturbative computations because one can factor the time and scale dependence of high-order diagrams (e.g., the contribution of order $n$ to the power spectrum scales as $D_+^{2n}$).  Beyond perturbation theory, several articles have recently been devoted to the study of exact ``consistency relations'' that remain valid in the nonperturbative regime, independently of the small-scale physics (including baryon or star-formation processes) \\cite{Kehagias:2013lr,Peloso:2013fk,Creminelli:2013qy,Kehagias:2013uq,Peloso:2013fj,Creminelli:2013kx,Valageas:2013ls,Creminelli:2013yq}. They relate the $(\\ell+n)$-point correlation, with $\\ell$ modes in the linear regime (soft domain) and $n$ modes at much higher wave numbers (hard domain) that can be in the nonlinear regime, to the $n$-point correlation (with $\\ell$ linear power spectrum prefactors). These results can be interpreted as the response of small structures (i.e., each element in the cosmic web such as walls, filaments or halos) to an initial density perturbation on much larger scales. More precisely, they derive from the equivalence principle, which ensures that all particles and structures fall in the same fashion in a gravitational potential force with a constant gradient. Then, at leading order, a large scale perturbation of the initial conditions merely transports smaller scale structures without distortions. Thus, a detection of a violation of these consistency relations would signal a deviation from Gaussian initial conditions, significant decaying modes, or a departure from general relativity.  In the standard scenario, the kinematic consistency relations discussed above vanish at equal times (because equal-time statistics cannot distinguish a uniform translation of the system). By going to the next order, and taking an angular average over the soft modes, Refs.~\\cite{Valageas:2013lr,Kehagias:2013fk} derived angular-averaged consistency relations that remain nontrivial even at equal times. Because this involves the dynamics of small-scale structures in a gravitational potential with a uniform curvature (the order beyond a constant gradient), this probes the physics beyond the equivalence principle and it is sensitive to the details of the dynamics. In particular, the explicit relations one obtains only hold for dark matter (i.e., they would be violated by nongravitational processes) and within the approximation $\\Omegam/f^2 \\simeq 1$, which enables us to relate the dynamics associated with different backgrounds (which correspond to different large-scale curvatures). However, within these approximations they remain valid in the nonperturbative regime.  In this study, we examine the validity of these angular-averaged relations by employing a large set of cosmological $N$-body simulations. We focus on the lowest-order consistency relation for the angular-averaged matter bispectrum, which is the most interesting one in practice. Reference~\\cite{Valageas:2013lr} has already checked this relation for the bispectrum explicitly at the leading order of perturbation theory. The aim of this study is to see how higher-order corrections enter both sides of the equation and how accurately the relation is recovered on smaller scales (i.e., whether and by how much nonlinearities amplify the inaccuracy due to the approximation $\\Omegam/f^2 \\simeq 1$).  This paper is organized as follows. We briefly review the angular-averaged consistency relations and their validation at tree level in Sec.~\\ref{sec:the_equation}. We then present the simulation analysis in Sec.~\\ref{sec:sim}, starting form the detail of the simulations in Sec.~\\ref{subsec:nbody} and next showing our results for the consistency relation in Sec.~\\ref{subsec:results}. We finally give a summary of the paper in Sec.~\\ref{sec:summary}. The effect of binning in the measurements of the spectra as well as a convergence study of the simulations are respectively presented in Appendixes~\\ref{app:binning} and \\ref{app:conv}. ",
        "conclusions": "One possible cause of the breakdown of the relation, which we find on small scales at $z=0.35$, is the approximation $\\Omegam/f^2 \\simeq 1$ employed in deriving the consistency relations. For instance, this combination deviates from unity by $\\sim7.5\\%$ at that redshift for the cosmological model we consider here, as already discussed in Fig.~\\ref{fig:om_o_f2}. We conduct some supplemental simulations to understand to what extent this affects the spectra of interest.    \\begin{figure} [ht!] % \\centering \\includegraphics[width=8.3cm]{EdS_pow.eps} \\caption{Ratio of the right-hand side of the consistency relation~(\\ref{eq:bispec}) in $\\Lambda$CDM and EdS background cosmology (triangles: $z=1$, circles: $z=0.35$). We also plot the contribution from each of the three terms by lines at $z=0.35$: solid, dashed, and dotted lines, respectively, show the first, second, and the last terms.} \\label{fig:EdS_pow} \\end{figure}    We run simulations with the same linear density field (i.e., the matter transfer function, the spectral tilt $n_\\mathrm{s}$, and the amplitude of the linear fluctuations $\\sigma_8$ at $z=0$) as the one we used in the main discussion, but adopting the cosmic expansion for the Einstein-de Sitter universe (i.e., $\\Omegam=1$; EdS hereafter, where we also have $\\Omegam/f^2=1$). We simulate four random realizations of $1024^3$ particles in a cubic box with $2048\\,h^{-1}\\mathrm{Mpc}$ a side with the same random seeds as the first four realizations of the low-resolution simulations in the main discussion.  Although these supplemental simulations are not self-consistent in a sense that they still adopt the transfer function calculated for a $\\Lambda$CDM model, they are useful to single out the effect of the breakdown of the approximation $\\Omegam/f^2 \\simeq 1$, by following the time evolution of the spectra starting with exactly the same initial values but in different backgrounds that probe different ranges of the ratio $\\Omegam/f^2$.  Note that, even though the EdS universe satisfies $\\Omegam/f^2=1$ at all times, at the background level, this is not sufficient to ensure the angular-averaged consistency relation is exact. Indeed, this relation assumes that $\\Omegam/f^2 \\simeq 1$ for all cosmologies close to the background one (as large-scale density fluctuations are identified with local changes of $\\Omegam$). Nevertheless, by changing the background cosmology, we change the reference point along the $\\Omegam/f^2$ ratio as a function of density fluctuations, and the comparison between the $\\Lambda$CDM and EdS cases gives an indirect probe of the effects associated with the approximation $\\Omegam/f^2 \\simeq 1$.  Since the function $D_+(z)$ is different between the two models, we choose the initial and output redshifts of the simulation in the EdS background such that they give the same linear growth rate: $D_{+,\\mathrm{EdS}}(z_\\mathrm{EdS}) = D_{+,\\Lambda\\mathrm{CDM}}(z_{\\Lambda\\mathrm{CDM}})$. If the approximation is accurate enough, in the sense that the nonlinear power spectrum only depends on the linear growth rate, these simulations should give spectra indistinguishable to those from the simulations in the main discussion. Also, we expect that the impact of any numerical error is almost the same in the two models since we are simulating exactly the same stage of structure formation with the choice of redshifts above.  We first show the ratio of the right-hand side of the relation~(\\ref{eq:bispec}) in Fig.~\\ref{fig:EdS_pow}. We plot the ratio of the sum of the three terms by triangles and circles, respectively, at $z=1$ and $z=0.35$. Although these symbols have vertical error bars estimated from the scatter among the four random realizations, they are hardly visible by eye: they are typically $10^{-5}$ to $10^{-4}$ level. This ensures the robustness of our estimate of the effect of $\\Omegam/f^2\\neq1$ using a rather small number of independent random realizations.  The plot shows that the effect is at most subpercent level at $z=1$, while a significant correction is observed at $z=0.35$. The correction is an increasing function of $k$ and reaches to $\\sim7.5\\%$ at $k=1\\,h\\,\\mathrm{Mpc}^{-1}$. We also show the ratio for each of the three terms separately by lines (solid, dashed, and dotted for the first, second, and the last terms, respectively). The solid line for the nonlinear matter power spectrum should be compared to similar analyses in the literature based on perturbation theories (e.g., Refs.~\\cite{Takahashi:2008lr,Pietroni:2008qy,Hiramatsu:2009fk}). Our simulation result is quantitatively in good agreement with these predictions in the literature on small scales (i.e., $k\\simlt 0.4\\,h\\,\\mathrm{Mpc}^{-1}$). The wave-number-derivative term, depicted by the dotted line, is also affected by the non-EdS background at a similar level as the solid line but toward the opposite direction. The largest effect lies in the time-derivative term (dashed).    \\begin{figure} [ht!] % \\centering \\includegraphics[width=8.3cm]{EdS_bis_z1.eps} \\caption{Ratio of the angular-averaged bispectrum [left-hand side of the consistency relation~(\\ref{eq:bispec})] in $\\Lambda$CDM and EdS background cosmology at $z=1$.} \\label{fig:EdS_bis_z1} \\end{figure} \\begin{figure} [ht!] % \\centering \\includegraphics[width=8.3cm]{EdS_bis_z035.eps} \\caption{Same as Fig.~\\ref{fig:EdS_bis_z1}, but at $z=0.35$.} \\label{fig:EdS_bis_z035} \\end{figure}     We now turn to the bispectrum and plot the ratio similar to the previous one in Figs.~\\ref{fig:EdS_bis_z1} and \\ref{fig:EdS_bis_z035}. We plot the ratio as a function of the soft wave number $k'$ for seven values of hard wave number $k$ shown in the figure legend. Again, the ratio is always close to unity at $z=1$, and the deviation is at most $\\sim 1\\%$ (at $k=1\\,h\\,\\mathrm{Mpc}^{-1}$). On the other hand, the ratio can be as large as $\\sim1.05$ at $z=0.35$ on small scales. The size of the deviation from unity increases with the hard wave number $k$, while its $k'$ dependence is weak at both redshifts.  An important observation here is that the correction from the non-EdS background affects both sides of the consistency relation~(\\ref{eq:bispec}), but its impact is quantitatively different in the two sides, which do not cancel out when we take their ratio. At $k=1\\,h\\,\\mathrm{Mpc}^{-1}$ at $z=0.35$, where we find that the consistency relation holds the worst, the net effect to the consistency relation is to lower the ratio of the left- to right-hand side by about $2\\%$ to $3\\%$ (see also Fig.~\\ref{fig:EdS_ratio} below). This qualitatively explains the ratio of the two sides smaller than unity in Fig.~\\ref{fig:ratio2}, though it does not completely recover the amplitude of the breakdown of the relation ($\\sim 7\\%$). However, this is not surprising because the comparison between the $\\Lambda$CDM and EdS cosmologies is only an indirect probe of the approximation $\\Omegam/f^2\\simeq 1$ (because neither of the two cosmologies is an exact reference point free from this approximation). These results suggest that a percent-level breakdown of the relation is naturally expected at $z=0.35$ on small scales for the $\\Lambda$CDM cosmology considered here at $z=0.35$.   \\begin{figure} [ht!] % \\centering \\includegraphics[width=8.3cm]{EdS_ratio2.eps} \\caption{Ratio of the two sides of the consistency relation~(\\ref{eq:bispec}) in either $\\Lambda$CDM (circles) or EdS (triangles) background cosmologies, as in Figs.~\\ref{fig:ratio1} and \\ref{fig:ratio2}. We focus on large values of $k$ at $z=1$ (left column) and $z=0.35$ (right column).} \\label{fig:EdS_ratio} \\end{figure}   We finally compare in Fig.~\\ref{fig:EdS_ratio} the consistency-relation ratios obtained in the $\\Lambda$CDM and EdS cosmologies, focusing on the nonlinear scales (we plot the data points from the four realizations for which the initial phases are reused in the EdS simulations to make the comparison fair). We obtain qualitatively similar results for both cosmologies. Moreover, in agreement with the previous figures, we find that for $k\\simlt 1\\,\\hmpci$ at $z=1$ and $k\\simlt 0.4\\,\\hmpci$ at $z=0.35$ the consistency relation is valid, whereas for $\\simgt 0.8\\,\\hmpci$ at $z=0.35$, the measured ratio is a few percent below unity, with a slightly smaller departure for the EdS case. As expected, the departure from unity takes place on scales where results from the EdS and $\\Lambda$CDM cosmologies start to deviate. To improve the theoretical predictions at low redshifts and small scales would require going beyond the approximation $\\Omegam/f^2 \\simeq 1$, but we do not investigate this problem here.     We have conducted a first numerical test of the angular-averaged consistency relation~(\\ref{eq:general-NL}) by exploiting a large suite of cosmological $N$-body simulations. We focus on the lowest-order example of the relation ($\\ell=1$ and $n=2$), which expresses the angular-averaged bispectrum in terms of the soft-mode and hard-mode power spectra. The large total volume of the simulations allows us to conduct a quantitative discussion on the validity of this relation.  We confirm that the relation is recovered within the statistical error of the simulations, beyond the validity range of the tree-level perturbation theory~(\\ref{eq:Btree}), for $k\\simlt 0.4\\,\\hmpci$ down to $z=0.35$. On the other hand, these scales remain within the range of higher-order perturbation theories so that the validity of the consistency relation is not surprising, because it is well known that the approximation $\\Omega_{\\rm m}/f^2$ used in most perturbative schemes is sufficiently accurate on these scales \\cite{Bernardeau02}. We indeed confirm that the breakdown of this approximation gives at most $\\sim 1\\%$ correction to the spectra, and the effect on the power and bispectra mostly compensate with each other.  Beyond this regime, we find that the validity range of the consistency relation extends to smaller scales, up to $k \\leq 1 \\,\\hmpci$ at $z=1$, where nonperturbative effects are not negligible \\cite{Valageas13a,Valageas2013}. We check that the condition $k'/k \\ll 1$ is not sufficient for the consistency relation, and the soft mode $\\tdelta_{\\bfk'}$ must be in the linear regime. Using the mixed bispectrum provides a more direct connection with the theory, and our results suggest that this also extends the validity range of the consistency relation. However, such a quantity can only be measured in numerical simulations and not from observations of the real Universe.  We find, on the other hand, a statistically significant breakdown of the relation at $z=0.35$, where the angular-averaged bispectrum is smaller than what we expect from the power spectrum. This feature is more prominent on smaller scales reaching to $\\sim7\\%$ at $k=1\\,\\hmpci$. We show that this is at least partly explained by the breakdown of the approximation $\\Omegam/f^2\\simeq 1$. We present an extensive convergence study of simulations in Appendix~\\ref{app:conv}, where we show that the systematic error to the ratio of the two sides is only mild if any. We thus conclude that the breakdown at this level should exist in reality in $\\Lambda$CDM cosmologies with parameters consistent with recent observations at low redshifts.  We leave further discussions on, for instance, the effect of nonlinear bias or the usefulness of the relation to detect primordial non-Gaussianity to future studies. Also, it might be interesting to see how baryonic effects alter the relation between different spectra in hydrodynamical simulations."
    },
    "1402/1402.5967_arXiv.txt": {
        "abstract": "{We investigate the connection between dust and gas in the nearby edge-on spiral galaxy NGC~891, a target of the Very Nearby Galaxies Survey. High resolution \\textit{Herschel} PACS and SPIRE 70, 100, 160, 250, 350, and 500~$\\mu$m images are combined with JCMT SCUBA 850~$\\mu$m observations to trace the far-infrared/submillimetre spectral energy distribution (SED). Maps of the \\hi \\ 21 cm line and CO(J=3-2) emission trace the atomic and molecular hydrogen gas, respectively. We fit one-component modified blackbody models to the integrated SED, finding a global dust mass of (8.5 $\\pm$ 2.0) $\\times$ 10$^{7}$ M$_{\\odot}$ and an average temperature of 23 $\\pm$ 2 K, consistent with results from previous far-infrared experiments. We also fit one-component modified blackbody models to pixel-by-pixel SEDs to produce maps of the dust mass and temperature. The dust mass distribution correlates with the total stellar population as traced by the 3.6~$\\mu$m emission. The derived dust temperature, which ranges from approximately 17 to 24 K, is found to correlate with the 24~$\\mu$m emission. Allowing the dust emissivity index to vary, we find an average value of $\\beta$ = 1.9 $\\pm$ 0.3. We confirm an inverse relation between the dust emissivity spectral index and dust temperature, but do not observe any variation of this relationship with vertical height from the mid-plane of the disk. A comparison of the dust properties with the gaseous components of the ISM reveals strong spatial correlations between the surface mass densities of dust (\\sdust) and the molecular hydrogen (\\shii) and total gas surface densities (\\sgas). These observations reveal the presence of regions of dense, cold dust that are coincident with peaks in the gas distribution and are associated with a molecular ring. Furthermore, the observed asymmetries in the dust temperature, the H$_{2}$-to-dust ratio and the total gas-to-dust ratio hint that an enhancement in the star formation rate may be the result of larger quantities of molecular gas available to fuel star formation in the NE compared to the SW. Whilst the asymmetry likely arises from dust obscuration due to the geometry of the line-of-sight projection of the spiral arms, we cannot exclude that there is also an enhancement in the star formation rate in the NE part of the disk.} ",
        "introduction": "Dust and gas are crucial ingredients for star formation in the Universe. Dust grains provide a surface on which molecular hydrogen (H$_{2}$) gas forms from atomic hydrogen (\\hi) gas via catalytic reactions (\\citealp*{gould1963}; \\citealp*{hollenbach1971}; \\citealp*{cazaux2002}). Regions of cooling and fragmentation of molecular gas in giant molecular clouds (GMCs) are the dominant formation sites of stars in galaxies (see \\citealp*{fukui2010}, and references therein). These stars produce the heavy elements via nucleosynthesis in their cores. Upon the death of a star, metals are expelled into the interstellar medium (ISM) and either mix with the gas phase or condense to form dust grains in the metal-rich cooling gas (see e.g., \\citealp*{nozawa2013}), such as in the expanding ejecta of novae and supernovae (e.g., \\citealp{Kozasa1991}; \\citealp{nozawa2003}; \\citealp{barlow2010}), the stellar winds from cool, dense atmospheres of asymptotic giant branch stars (e.g., \\citealp*{ferrarotti2006}; \\citealp*{zhukovska2013}), and the common-envelope in close binary systems (e.g., \\citealp{lu2013}). Growth, modification and destruction processes may affect the overall abundance of interstellar dust (e.g., \\citealp{clemens2010}; \\citealp{mattsson2012}), whilst surviving grains have the potential to aid future episodes of star formation. Thus, the evolutionary processes governing gas, stars, metals and dust are mutually related, and constraining the relationships between these four components is indispensable for understanding galaxy formation and evolution.  Detailed studies of the interaction between gas and stars have been technologically feasible for a long time. Global star formation rates (SFRs) of galaxies are observed to be well correlated with their atomic gas content (e.g., \\citealp{schmidt1959}; \\citealp{kennicutt1998}). There is a power-law correlation between the surface densities of gas (\\sgas) and SFR (\\ssfr), with numerous studies attempting to constrain the power-law index via global quantities from galaxy samples and resolved observations of individual systems (see \\citealp*{kennicutt2012} for a review). Spatially resolved data has uncovered that \\ssfr \\ is more strongly correlated with the surface density of molecular hydrogen (\\shii) than atomic hydrogen (\\shi) gas \\citep*{wong2002}. More recent studies have begun to probe this relationship at sub-kpc scales, approaching the typical size of GMCs and the intrinsic physical scale of star formation, for larger samples of galaxies (e.g., \\citealp{bigiel2008}; \\citealp{momose2013}). Strong observational evidence also shows a correlation between the stellar mass and the gas-phase metallicity (e.g., \\citealp{lequeux1979}; \\citealp{tremonti2004}). Recent efforts are attempting to observationally constrain the relationships between stellar mass, metal content and \\hi \\ gas content (e.g., \\citealp{hughes2013}; \\citealp{bothwell2013}).  Studying the dust content has proved a difficult endeavour mainly due to technological limitations. Dust absorbs up to 50\\% of the optical and ultraviolet (UV) photons from stars. The dust can then be traced via the re-emission of this energy at mid-infrared (MIR), far-infrared (FIR) and submillimetre (submm) wavelengths. Previous space missions - necessary to avoid atmospheric absorption in these bands - such as IRAS, ISO, \\textit{Spitzer} and AKARI (respectively: \\citealp{neugebauer1984}; \\citealp{kessler1996}; \\citealp{werner2004}; \\citealp{murakami2007}) typically operated at wavelengths  $\\lambda \\sim$~240~$\\mu$m, potentially missing a great part of the spectral energy distribution (SED) associated with a cold dust component. Fortunately, the success of the \\textit{Herschel} Space Observatory \\citep{pilbratt2010} satellite has substantially improved this situation. With two of its instruments, the Photodetector Array Camera and Spectrometer \\citep[PACS,][]{poglitsch2010} and the Spectral and Photometric Imaging REceiver \\citep[SPIRE,][]{griffin2010}, \\textit{Herschel} observations are capable of constraining the FIR SED, and thus detect emission at $\\lambda >$~250~$\\mu$m originating from a cold dust component (e.g., \\citealp{bendo2010b}), via imaging in six wavebands from 70 to 500~$\\mu$m with significantly higher sensitivity and angular resolution than previous FIR/submm experiments.  Recent observational studies have used \\textit{Herschel} data to investigate the correlations between the dust component and other galaxy properties (see e.g., \\citealp{smith2012a}; \\citealp{bourne2013}; \\citealp{remyruyer2013}). For example, scaling relations have been found between the total masses of dust, H$_{2}$ and \\hi \\ in late-type Virgo cluster galaxies \\citep{corbelli2012}, between the dust-to-stellar mass ratio and $NUV-r$ colour (\\citealp{cortese2012}), and between the dust temperature and SFR per unit dust mass \\citep{clemens2013}. A growing number of studies are using the unprecedented angular resolution of the \\textit{Herschel} observations to perform pixel-by-pixel fitting of the FIR/submm SED and map the main dust properties of nearby galaxies (e.g., \\citealp{smith2010}; \\citealp{galametz2012}; \\citealp{boquien2013}; \\citealp{draine2013}; \\citealp{tabatabaei2013}). Observations obtained by the \\textit{Herschel} Exploitation of Local Galaxy Andromeda \\citep{fritz2012} enabled highly-detailed maps of the dust properties of M31 to be produced via such pixel-by-pixel SED fitting \\citep{smith2012}, and allow for an in-depth study of Andromeda's ISM (\\citealp{ford2013}; \\citealp[][in prep.]{viaene2013}; see also \\citealp{draine2013}). Similarly, \\citet{foyle2012} and \\citet{mentuchcooper2012} have examined the gas and dust using \\textit{Herschel} observations to carry out a pixel-by-pixel analysis of the inner regions of M83 and M51, respectively. Neither galaxy shows strong variations in the dust-to-gas mass ratio, except for a small gradient in M51 \\citep{mentuchcooper2012}, whereas many other galaxy studies have found strong gas-to-dust ratio gradients (e.g., \\citealp{munozmateos2009}; \\citealp{bendo2010}; \\citealp{magrini2011}). Although much effort has been made to quantify the radial dust distribution in these face-on spiral galaxies, we also require information on the vertical $|z|$ structure for a complete picture of the three dimensional properties of dust and gas in galaxies.  Edge-on spiral galaxies with inclinations at or near 90$^{\\circ}$ offer the best opportunity to study the radial and vertical structure of the various galaxy components. The stellar component has been shown to be sensitive to the merging history of a galaxy, whereas the ISM may probe gravitational instability (see e.g., \\citealp{dalcanton2004}; \\citealp*{yoachim2006}). Dust lanes feature prominently in optical images of massive edge-on spirals, enabling the study of dust in both emission and absorption. Furthermore, the increased surface brightness arising from the line-of-sight projection allows for both the vertical and radial dust distribution to be studied in detail at greater distances from the centre. Edge-ons therefore provide an important contribution to our knowledge of the dust structure. Fortunately, we have one such prototypical example of an almost perfect edge-on spiral right in our neighbourhood: NGC~891. Located at a distance of $\\sim$9.6 Mpc (e.g., \\citealp{strickland2004}) with an inclination angle of $>$89$^{\\circ}$  (e.g., \\citealp{xilouris1998}), this bright, non-interacting SA(s)b galaxy \\citep{RC21976} presents similarities to our own Milky Way in Hubble type, rotational velocity and optical luminosity (see Table~\\ref{tab:basicproperties}). These properties make NGC~891 an ideal target for studying the stars, gas and dust in the disk, and so it has already been extensively observed at a range of wavelengths.  The stellar content in the galaxy, traced in the optical wavelengths (e.g., \\citealp*{vanderkruit1981}; \\citealp{ibata2009}) and at 3.6~$\\mu$m \\citep{kamphuis2007}, are found to be fairly symmetric along the galactic plane. However, the \\hi \\ gas shows a slight north-south asymmetry in both disk and halo thickness \\citep{swaters1997}, and extends up to at least 22 kpc above the galactic plane \\citep{oosterloo2007}. In contrast, molecular hydrogen gas observed via CO emission is predominantly found in a thin disk around 400 pc thick \\citep{scoville1993}, and also seen in the halo up to 1.4 kpc above the plane \\citep{gb1992}. The H$\\alpha$ emission, a tracer of ionized gas linked to star formation, is more prominent and extended in the $|z|$ direction on the northern side of the galaxy than the southern (\\citealp{dettmar1990}; \\citealp{rand1990}). Whether the north-south asymmetry arises due to higher SFR in the north \\citep{rossa2004} or increased attenuation by dust in spiral arms along the line-of-sight through the disk (\\citealp{kamphuis2007}; \\citealp{schechtmanrook2012}) remains an open issue, which may be resolved by studying the cold dust in the galaxy.  Evidence for cold dust in NGC~891 first came from IRAS observations at FIR wavelengths (\\citealp{wainscoat1987}), and later at millimetre wavelengths using the IRAM 30~m telescope (\\citealp{guelin1993}). The 1.3~mm dust emission was noted to correlate strongly with the H$_{2}$ gas traced by CO emission, but poorly with the \\hi \\ emission. Later observations at $\\lambda\\,$ 450 and 850~$\\mu$m with the Submillimeter Common-User Bolometer Array (SCUBA) at the James Clerk Maxwell Telescope (JCMT) found a large amount of cold dust ($\\sim 15$\\,K) in the disk \\citep{alton1998}. The submm emission was confirmed to correlate spatially with the molecular gas in the disk (\\citealp{alton1998}; \\citealp{israel1999}). The combination of the SCUBA observations, IRAS and ISO data, with maps from the French balloon-borne PRONAOS experiment \\citep{serra2002} traced the complete FIR/submm SED of NGC~891 for the first time \\citep{dupac2003a}, finding a dust temperature of $\\sim$ 18 - 24 K and estimating a global gas-to-dust mass ratio of $\\sim$ 240. \\citet{popescu2004} present deep maps of dust emission at  170 and 200~$\\mu$m obtained with the ISOPHOT instrument. However, the low resolution of these datasets, typically $\\geq$ 1\\arcmin , meant a pixel-by-pixel analysis of the FIR/submm SED to map the cold dust distribution was not previously feasible.  New, high resolution ($\\leq$ $36\\farcs4$) \\textit{Herschel} PACS/SPIRE observations of NGC~891 have been obtained as part of the \\textit{Herschel} Guaranteed Time Key Project the Very Nearby Galaxies Survey (P.~I.:~C.~D.~Wilson), which aims to study the gas and dust in the ISM of a diverse sample of 13 nearby galaxies using \\textit{Herschel}. These observations now make it possible to trace the FIR/submm SED in detail and map the resolved dust properties in the galaxy disk. In this paper, we investigate the connection between dust and gas in NGC~891 in order to test for asymmetries in the distributions of the dust properties and the gas-to-dust ratio. We use these new \\textit{Herschel} PACS/SPIRE observations combined with a JCMT SCUBA 850~$\\mu$m image to trace the FIR/submm SED. For the first time, we fit one-component modified blackbody models to the SEDs of each pixel to map the dust mass and temperature in the galaxy. HI 21 cm line emission and $^{12}$CO(J=1-0) emission maps, combined with a new JCMT HARP-B $^{12}$CO(J=3-2) map (P.~I.:~T.~Parkin), trace the atomic and molecular hydrogen gas, enabling the study of the gas-to-dust ratios of these components. Our paper is structured as follows. In the next section, we present the new and existing data used in our analysis. In Section 3, we present our detection of extraplanar dust. In Section \\ref{sec:dustprops}, we describe our SED fitting methodology and the resulting integrated and resolved dust properties. Section \\ref{sec:dustgas} describes the relationships we find between the dust and gas, and Sections \\ref{sec:discussion} and \\ref{sec:conclusions} present our discussion and conclusions, respectively.   \\begin{table} \\centering \\begin{minipage}{\\columnwidth} \\caption{Basic properties of NGC 891} \\label{tab:basicproperties} \\begin{tabular}{lll} \\hline \\hline Property & value & reference\\\\ \\hline RA (J2000) & $2^\\mathrm{h}\\ 22^\\mathrm{m}\\ 33\\fs0$ & NED\\footnote{The NASA Extragalactic Database (NED) is available online at \\\\ \\url{http://nedwww.ipac.caltech.edu/}} \\\\ Dec (J2000) & +$42\\degr\\ 20\\arcmin\\ 57\\farcs2$ & NED \\\\ Hubble type & SA(s)b & \\citet{RC21976} \\\\ distance & 9.6 Mpc & Strickland et al. (2004) \\\\ optical major axis & 13.5 arcmin & \\citet{RC31991} \\\\ optical minor axis & 2.5 arcmin & \\citet{RC31991} \\\\ major axis angle & $22^\\circ$ & NED \\\\ inclination angle & $89^\\circ$ & \\citet{xilouris1998} \\\\ B-band luminosity & 7.8 $\\times$ $10^{9}$ $\\mathrm{L}_\\odot$ & \\citet{RC31991} \\\\ rotational velocity & 225 km s$^{-1}$ & \\citet{vanderkruit1984} \\\\ total mass & 1.4 $\\times$ $10^{11}$ $\\mathrm{M}_\\odot$ & \\citet{oosterloo2007} \\\\ SFR & 3.8 $\\mathrm{M}_\\odot$ yr$^{-1}$ & \\citet{popescu2004}\\\\    % \\hline \\end{tabular} \\end{minipage} \\end{table}  \\newpage ",
        "conclusions": "\\label{sec:conclusions}  In this work, we investigated the connection between dust and gas in the nearby spiral galaxy NGC~891. High resolution \\textit{Herschel} PACS and SPIRE 70, 100, 160, 250, 350, and 500~$\\mu$m images are combined with SCUBA 850~$\\mu$m observations to trace the far-infrared/submillimetre SED. We fit one-component modified blackbody models to the integrated SED, finding a global dust mass of (8.5 $\\pm$ 2.0) $\\times$ 10$^{7}$ M$_{\\odot}$ and temperature of 23 $\\pm$ 2 K, consistent with results from previous far-infrared experiments. We also fit one-component modified blackbody models to pixel-by-pixel SEDs to produce maps of the dust mass and temperature at physical scales of $\\sim$ 0.5 kpc. The dust mass distribution shows evidence of several peaks along the disk. The derived dust temperature, which ranges from approximately 17 to 24~K, does not correlate with the dust spatial distribution. Allowing the dust emissivity index to vary, we find an average value of $\\beta$ = 1.9 $\\pm$ 0.3. We confirm an inverse relation between the dust emissivity spectral index and dust temperature, but do not observe any variation of this relationship with vertical height from the mid-plane of the disk. We adopt a fixed value of $\\beta$ = 1.8 for our main analysis comparing the derived dust properties to the gaseous components.  Combining the mapped dust properties with observations of the \\hi \\ 21~cm line emission and CO(J=3-2) emission to trace the atomic and molecular hydrogen gas, respectively, we examined the correlations between the dust properties with the gaseous components of the ISM. We found strong spatial correlations between the surface mass densities of dust (\\sdust) and the molecular hydrogen (\\shii) and total gas surface densities (\\sgas). We observed no correlation between the dust content and the \\hi \\ gas content. These observations reveal the presence of regions of dense, cold dust that are coincident with peaks in the gas distribution and appear associated with a molecular ring. Furthermore, the observed asymmetries in the dust temperature, the H$_{2}$-to-dust ratio and the total gas-to-dust ratio hint that an enhancement in the star formation rate may be the result of larger quantities of molecular gas available to fuel star formation in the NE compared to the SW. Whilst we suspect that the asymmetry likely arises from dust obscuration due to the geometry of the line-of-sight projection of the spiral arms, we cannot exclude that there is also an enhancement in the star formation rate in the NE part of the disk.  This study raises several intriguing questions. The observed relationship between the dust temperature and the surface density of the molecular hydrogen suggests that dust may be cooler in dense regions of dust and molecular hydrogen, perhaps due to the processes of either self-shielding or molecular gas shielding from heating by the UV interstellar radiation field, or dust coagulation. However, whilst these processes have support from observations at high spatial resolution in nearby molecular clouds, it is unclear whether our lower resolution dataset is capable of uncovering such a correlation between the dust temperature and molecular gas density observed on these large spatial scales. This remains an open issue, which may be resolved by a detailed analysis and modelling of the $T_{\\mathrm{dust}}$-\\shii \\ relation based on larger samples of edge-on galaxies, e.g. \\textit{HERschel} Observations of Edge-on Spirals \\citep[\\textit{HER}OES, ][]{verstappen2013} and the New \\textit{HErschel} Multi-wavelength Extragalactic Survey of Edge-on Spirals \\citep[NHEMESES, ][]{holwerda2012}. Furthermore, it is still an open question whether molecule formation is mainly driven by hydrostatic pressure or UV radiation shielding over different spatial scales and over a large range of metallicities (see e.g., \\citealp{fumagalli2010} and references therein). Upcoming facilities such as the Atacama Large Millimeter Array and the Square Kilometer Array should shed some light on this problem. These interferometers will allow the mapping of atomic and molecular gas and dust at high resolution, potentially enabling the study of how the relative quantities of these ISM components vary with metallicity and photodissociating UV radiation intensity in different environments. A formalisation of the processes governing the molecular content in galaxies will undoubtedly aid future cosmological simulations of galaxy formation and evolution.  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{dust_iracmips.eps} \\end{center} \\vspace{-0.3cm} \\caption[Dust properties versus IRAC and MIPS fluxes]{The pixel-by-pixel relationships between the derived dust properties, \\sdust \\ (\\textit{upper panels}) and $T_{\\mathrm{dust}}$ (\\textit{lower panels}), and the IRAC~3.6~$\\mu$m (\\textit{left panels}) and MIPS~24~$\\mu$m flux densities (\\textit{right panels}). The best linear fits are indicated by the red solid lines. Reported $\\rho$ values are the corresponding Pearson correlation coefficients.}\\label{fig:dustsourcecomps} \\end{figure}"
    },
    "1402/1402.0403_arXiv.txt": {
        "abstract": "{We analyze  60 months of all sky data from the {\\it Fermi}-LAT. The Fermi Bubble structures  discovered previously are clearly revealed by our analysis. With more data and, consequently, better statistics we can now divide each bubble into  constant longitude slices to investigate their gross $\\gamma$-ray spectral morphology. While the detailed spectral behaviour of each slice derived within our analysis is somewhat dependent on the assumed background model, we find, robustly,  a relative deficit of the flux at low energies (i.e., hardening) towards the top of the South Bubble. In neither Bubble does the spectrum soften with longitude. The morphology of the Fermi Bubbles is also revealed to be energy dependent: at high energies they  are more extended. We conclude from the gamma-ray spectrum at high latitudes that a low energy break in the parent cosmic ray population is required in both leptonic and hadronic models. We briefly discuss  possible leptonic and hadronic interpretation of this phenomenology. } ",
        "introduction": "Two huge, bubble-like structures have been reported by \\citet{su10}, \\citet{dobler10}, and \\citet{su12} in {\\it Fermi}-LAT gamma ray data to extend $\\sim$50$^\\circ$ above and below the Galactic center. The gamma ray emission from these structures, dubbed the `Fermi Bubbles' (FBs), exhibits a $ E^{-2}$ type spectrum, significantly harder than  the spectrum of the diffuse gamma-ray emission from the Galactic disk. Remarkably, structures coincident or similar to the FBs can be seen at other wavelengths, including the  (total intensity) microwave haze found in WMAP  \\citep{finkbeiner04} and, most recently, in Planck data \\citep{planckhaze}; the polarized microwave structures reported by \\citet{Jones2012}; the large-scale, biconical structures found \\citep{Bland-Hawthorn2003} in ROSAT X-ray data \\citep{rosat}; and the giant polarised radio lobes recently found at 2.3 GHz \\citep{carretti13}. Several  models  have been proposed to explain both the morphology and spectral properties of the FBs \\citep{cheng11, mertsch11, crocker11, crocker12, zubovas11, zubovas12, guo2011, guo2012, yang2012,Crocker2013}. These models predict different energy dependent morphologies. For example, in the simplest IC model  (e.g., discussed by Su et al.~2010) with a spatially-independent electron spectrum described by  a power law electron population with a low energy cutoff at $500~\\rm$ GeV, one expects   a tendency for softening of the gamma-ray spectra at high latitudes. This is due to   the gradual reduction of the  IC  components  produced via  upscattering of the optical and UV photons  - the main target fields contributing to the  production of gamma-rays above $10~\\rm$ GeV that also decline towards high latitudes. On the other hand, in a simple one-zone hadronic model \\citep[e.g.,][]{crocker11} the protons' steady state distribution should produce a similar spectrum of gamma rays at all latitudes. Thus, studies of the energy-dependent morphology of the FBs may shed light on the nature of the gamma ray emission mechanism(s). This is the basic motivation for the current study.   Here we present an analysis based on 60 months' $Fermi$-LAT data. We find a similar overall morphology and spectrum for the FBs  to those obtained by \\citet{su10}. The FBs exhibit a rather homogeneous {\\it surface} brightness. Importantly, this implies  non-homogeneous $volumetric$ emissivities considering projection effects \\citep{su10}. The FBs, however, are not completely uniform: they exhibit some `hot spots' (referred to by \\citet{su10,su12} as the `jet', `donut' and `cocoon').   In this paper the significantly larger photon statistics and the availability of the recently updated  Fermi science software tools \\citep{pass7} allow us to investigate the gross spectral morphology of the FBs. A somewhat similar study has recently been conducted by \\citet{Hooper2013}. However, they concentrate spectral features in low latitudes while we focus on the  high latitude, especially the top of the bubble. We proceed by dividing each bubble into several slices to investigate possible gamma-ray spectral change with latitude. Employing different background models, we find, robustly, a spectral hardening, or more specifically,  a deficit of low energy flux  towards the top of the SFB. We are also able to extend the  spectrum of the FBs to lower energy than previously attempted because of the improvement of the analysis software and the instruments response functions. The paper is structured as follows: in section 2 we  present the results of our data analysis, in section 3 we discuss the fitting of leptonic and hadronic models to the gamma ray data, and in section 4 we set out our conclusions. ",
        "conclusions": "The gamma-ray residual maps obtained in the current study -- based on six years' {\\it Fermi}-LAT data -- confirm that the surface brightness of  each bubble is  homogeneous at a gross level. This implies that the {\\it volumetric} emissivity is not homogeneous, otherwise we would detect a higher brightness in the center of each bubble due to projection \\citep{su10}. One possible explanation of this may be higher turbulence in the bubbles' edges generating more efficient particle acceleration or stronger confinement of high energy particles there with a resulting increase in the local gamma ray emissivity; alternatively, in a non-saturation hadronic model, higher gas densities towards the edges may have the same effect \\citep{Crocker2013}.   Below we assume the low energy deficit  of the gamma ray spectrum at high latitudes found above is a real  effect and discuss its possible implications.  \\\\   To simplify our modeling, we assume each bubble is  a sphere with a radius 3.4 kpc. The 2-d projection of each bubble can be approximated as a half sphere in the region $b>30^{\\circ}$ for the NFB and $b<-30^{\\circ}$ for the SFB. The center of each bubble is located at (0,0,$\\pm$5kpc) in  Galactic coordinates. Assuming an {\\it intrinsic} north-south symmetry, we only analyse the SFB and, to reveal most starkly the spectral change with latitude, we consider only slice South 1 ($-33^\\circ <b <-25^ \\circ $) and South 4 ($-55^\\circ <b <-47^ \\circ $) in detail. Slice South 1 covers the center of the SFB. We use the ISRF value in Galprop at  (0,0,5 kpc)  for the average value of the slice South 1 in the calculation, $w_{IR}= 0.18$ eV/cm$^3$ and $w_{opt}= $ 0.9 eV/cm$^3$. The height of South 4 is about 7 kpc, thus we adopt the ISRF energy density from Galprop at (0,0,7 kpc),  $0.7 $ eV/cm$^3$ for the optical component and $0.15 $ eV/cm$^3$ for the IR component. The IR  and optical photon fields are modeled as diluted blackbody (gray body) spectra with temperatures of 100 K and 5000 K, respectively.\\\\    In the hadronic scenario, motivated by the relative low energy deficit of the gamma ray flux in the top slice, we fit the SED in South 1 with a pure power law proton spectrum while for South 4 we use a power law with a low energy break that we find needs to be at $20~\\rm GeV$ to fit the data. We assume the proton flux below the  break is zero for simplicity (see Fig.~\\ref{proton_spec}). Our results are shown in Fig.~\\ref{fig:pp}. The required low energy break may arise naturally due to   energy-dependent diffusion effects in a non-steady-state scenario: given South 4 is far from the (assumed) injection source in the plane, it may be that only high energy particles (which diffuse faster) have had time to travel there since a previous injection event or given the age of the structure. The formalism describing these energy-dependent diffusion effects can be found in \\citet{aa96} where both  impulsive and continuous injection cases are described. The position of the low energy break can be estimated by equating the age of the bubble (or the time since a previous burst) with the diffusion time, i.e.,the beak energy is  $E_{\\rm bk} \\simeq (d^2/(D_{\\rm 0} t ))^\\frac{1}{\\delta}$, where $d$ is the distance of the top slice from the injection source, $t$ is the age of the bubble (or burst) and $D(E)=D_{\\rm 0}E^{\\delta}$ is the energy dependent diffusion coefficient. If we assume  protons are injected at the Galactic center, $d$ is about $7~\\rm kpc$. Notice that in the kinetic equations which govern  proton diffusion only the combination $D_{\\rm 0} \\  t$ appears. Thus any timescale can be obtained if we tune the diffusion coefficient $D_{\\rm 0}$. If $D(E)$ takes a value similar to that in the Galactic plane, say  $4\\times10^{28} $ cm$^2 $ s$^{-1}$ at 1 GeV and $\\delta=0.3$, the age of the FBs should be $t \\sim  10^8$ yr (cf.~\\citealt{Crocker2013}). Alternatively, if there is continuous proton injection into the FBs over $10^{10}$~yrs we need a rather small diffusion coefficient, $D_0 \\sim 10^{26}$ cm$^2$ s$^{-1}$ at 1 GeV, which is about 1/100 of that in the Galactic plane.   If we assume the injection is impulsive and  the ISM density is  $0.005$ cm$^{-3}$ the total energy in protons needed to light the Bubbles is of the order $10^{56} ~\\rm erg$. Note that if we assume the simple geometry mentioned above and a magnetic field $B \\sim 10~\\rm \\mu G$ as suggested by \\citet{carretti13}, then the energy densities of cosmic rays and magnetic fields are close to equipartition. A burst-like  injection event requires that the duration of the injection is much smaller than the age of the structure which is $t \\sim  10^8$ yr (for $D_{\\rm 0} = 4\\times10^{28} $ cm$^2 $ s$^{-1}$ at 1 GeV), i.e., the duration of the injection event would have to be of order $10^7 ~\\rm$ yr or less in this case. Then the injection rate is of the order $10^{41}$ erg/s, two orders of magnitude larger than the the X-ray luminosity of a X-ray reflection nebulae near the Galactic center \\citep[e.g.,][]{ryu12}, but still only one thousandth of the Eddington Luminosity of Sgr $A^*$, $\\sim 5\\times 10^{44}~\\rm erg/s$.    \\begin{figure*} \\centering \\includegraphics[width=70mm,angle=0]{proton_spec.eps} \\includegraphics[width=60mm,angle=0]{elec_spec.eps} \\caption{Assumed CR proton (left) and electron (right) distributions in Slices 1 and 4 of the SFB}. \\label{proton_spec} \\end{figure*}       A possible concern for an energy-dependent diffusion scenario is that, in the high energy range for which the diffusion length is much larger than the distance to the source, the predicted proton density should be the same at all latitudes for the impulsive case, or scale as $1/R$ (where $R$ is the distance to the injection source) for the continuous case. Thus, given the (expected) relatively smaller line of sight through the top of the SFB, one naively anticipates that the gamma ray flux  in South 4 is smaller than that in South 1. The observed data reveal a similar total flux in both slices at high energy, however. One possible resolution to this tension may be that the target gas density is higher at high latitudes. It is also likely that the geometry of the FBs is different from our simple  assumption of sphericity thereby entailing non-trivial projection effects.  \\\\  In a leptonic scenario the low energy deficit of gamma ray flux in the top slice also requires a corresponding low energy break in  the electron distribution (see Fig.~\\ref{proton_spec} right). This  can, again, be produced by several mechanisms. Firstly, the energy-dependent diffusion effects mentioned above also work for electrons in principle \\citep[cf.][]{McQuinn2011}. However the much faster cooling of electrons relative to protons tends to prevent such effects playing an important role. To fit the high energy gamma ray data we need TeV electrons, whose cooling time scale is about $10^5~\\rm yrs$ in the circumstances of the FBs. Thus, even if we assume very fast diffusion with the diffusion coefficient of $4\\times 10^{29}$ cm$^2$/s at 1 GeV (and an index of 0.6), the diffusion distance of TeV electrons in $10^5~\\rm$ yrs is only 2 kpc, significantly less than the distance of South 4 from the plane (which is more than 5 kpc). This implies high energy electrons accelerated at low latitudes can never reach South 4 and, as a result, if we want to explain the radiation of the FBs in a leptonic scenario, the electrons should be accelerated {\\it in situ}. \\\\  The second problem related to electrons concerns  the interpretation of the very hard energy spectrum of radiation below $5~\\rm GeV$.  Such a hard spectrum formally can be explained by assuming a low energy break in the electron spectrum.    However,  if the electrons are cooled,  no matter how hard the injection spectrum is, because of the radiative  ($dE/dt \\sim E^2$ type) losses the resulted {\\it steady-state} distribution below the break in the injection spectrum,  will have have a standard, $E^{-2}$ power law-spectrum. The latter will result in IC gamma-ray spectrum  with a photon index 1.5 which  is still not sufficient to explain the observed gamma-ray spectrum  of South~4; below  $5~\\rm GeV$ the photon index of the latter   is as small as $1$.  One can formally  overcome this problem assuming that the electrons stay uncooled in FBs.  However in order to reproduce the break in the gamma-ray spectrum, one would need a corresponding break at $100~\\rm GeV$. The corresponding synchrotron cooling time at this energy,  $t \\sim 8 \\times 10^5 (\\frac{100~ \\rm GeV}{E})(\\frac{B}{10~ \\rm \\mu G})^{-2} ~\\rm yr$,  appears  significantly shorter than  the age of  FBs within any  reasonable  model  of the latter. Thus we  need very efficient continuous acceleration acceleration of electrons which would dominate over the rate of energy losses and in that way keep very hard the steady-state spectrum. This in principle can be realized  with stochastic acceleration which can produce Maxwellian  type {\\it steady-state} energy distribution of electrons. Indeed adopting  an electron distribution like $N(E) \\sim E^2 \\ exp(-E/E_0)$ with  $E_0= 300~\\rm$ GeV,  we can have a reasonable fit to the detected gamma-ray spectrum  in South 4, as it is shown in Fig.~\\ref{fig:pp}.  For comparison, we also  show  IC emission at the position of South 1 and South 4 by adopting a  broken power law electron distribution with an index $2$ below the break at $200~\\rm GeV$ and $3$ above.  The results fit the  SED of South 1 well, but fail to fit that of South 4, which are also shown in Fig.~\\ref{fig:pp}.            Anisotropy effects may, in principle, increase the IC  gamma ray flux at high latitudes since most  of  seed photons come from regions  in  the Galactic plane;   thus gamma-rays we detect at high latitudes are  a result of (almost)  head-on collisions. To probe  this effect we use the formalism described in \\citet{aniic}. The enhancement factor is defined by the flux ratio of the anisotropic IC  to the isotropic IC  assuming the the same ISRF energy density. Results for the positions of South 1 and South 4 are shown in Fig.~\\ref{fig:pp}. We can see that although the anisotropic IC effect may  cause a spectral a hardening, however  the difference between South~1 and South~4 is rather  small, less than $10$\\%. It appears not sufficient to explain the  very hard spectrum of South~4.           \\begin{figure*} \\centering \\includegraphics[width=50mm,angle=0]{pp_pl.eps} \\includegraphics[width=50mm,angle=0]{icfit.eps} \\includegraphics[width=50mm,angle=0]{cf.eps}  \\caption{ Left and Middle panel:  Fitting the FBs' SED with the hadronic model (left panel)  and the leptonic model (middle panel) described in the text. The red  curve shows the fitted spectrum and the red hatched region spans  all spectra resulting with from different templates for South 1;  blue curve and hatched region are for South 4. For leptonic models,  for South 1 we assume the broken power law described in the text and the ISRF value given by GALPROP. For South 4, the black curve is for the same broken power law with the ISRF floating while for the gray curve we use the ISRF in GALPROP at the position of South 4 and a Maxwellian type electron distribution with $E_{\\rm bk}=300~\\rm GeV$. In our hadronic scenario  energy dependent diffusion effects are considered as  described in the text. Right panel:  The enhancement factor due to anisotropic IC. The IR and Opt/UV components of the ISRF have effective  temperatures of $100$  K and $5000$ K, respectively. Results for $z=5~\\rm kpc$ and $z=7~\\rm kpc$ are shown (approximate heights of South 1 and South 4 above the Galactic plane).   } \\label{fig:pp} \\end{figure*}"
    },
    "1402/1402.5004_arXiv.txt": {
        "abstract": "In this paper we study the creation of super-massive real scalar particles in the framework of a $f(R)=R-\\beta/R^n$ modified gravity theory, with parameters constrained by observational data. The analysis is restrict to a homogeneous and isotropic flat and radiation dominated universe. We compare the results to the standard Einstein gravity with cosmological constant ($\\Lambda CDM$ model), and we show that the total number density of created particles in the $f(R)$ model is very close to the standard case. Another interesting result is that the spectrum of created particles is $\\beta$ independent at early times. ",
        "introduction": "Recently, non-standard gravity theories have been proposed as an alternative to understand the physical mechanism behind the late-time acceleration of the universe. This is naturally obtained in $f(R)$ gravity theories (for a review, see \\cite{sotiri}). In such an approach, the curvature scalar  $R$ in the Einstein-Hilbert action is replaced by a general function $f(R)$, so that the standard Einstein field equation is recovered as a particular case \\cite{fR,allemandi,allemandi2,sotiri2,meng,dolgov,cembranos,gaspe,vilk}. A review of various modified gravities considered as gravitational alternative for dark energy, an unified description of the early-time inflation and the solutions of some related problems in $f(R)$ theories have been presented by Nojiri and Odintsov \\cite{nojiri}.  Cosmological constraints on $f(R)$ gravity theories has also been studied recently \\cite{carvalho,elgaroy,fay,koivisto}. Of particular interest is the model $f(R)=R-\\beta/R^n$, investigated by Carvalho et al. 2008 \\cite{carvalho}. They used determinations of the Hubble function, H(z), based on differential age method to place bounds on the free parameters $n$ and $\\beta$, combined with constraints from Baryon Acoustic Oscillations (BAO) and Cosmic Microwave Background (CMB) measurements. The best-fit values for $n$ and $\\beta$ obtained from H(z)+BAO+CMB test are $0.03$ and $4.7$, respectively, with the parameters $n$ to lie in the intervals $n \\in [-0.25,0.35]$ and $\\beta \\in [2.3,7.1]$ at 99.7\\% c.l. The standard $\\Lambda CDM$ model corresponds to $n=0$ and $\\beta = 4.38$.  By considering a radiation dominated universe, a recent work \\cite{saulocarlos} studied the quantum process of particle creation in a $f(R)=R+\\beta R^n$ theory. Although very similar to the $f(R)$ considered by Carvalho et al. 2008 \\cite{carvalho}, the authors analyzed  only  positive $\\beta$ values and for some specific values of $n$, namely $n= (2+2j)/(2j+1)$ and $n= 2j/(2j+1)$, $j=1,2,3,\\dots $. It was shown that both massive and massless scalar particles can be produced by purely expanding effects.  In the present work we investigate the particle production in the model $f(R) = R -\\beta/ R^n$, with the  parameters obtained by Carvalho et al. 2008 \\cite{carvalho}, $n \\in [-0.25,0.35]$ and $\\beta \\in [2.3,7.1]$ at 99.7\\% c.l., in a homogeneous and isotropic flat and radiation dominated universe. We show that super-massive particles can be created at the early times, when the universe is radiation dominated. This could be an alternative mechanism to create dark matter particles in the course of the evolution of the universe. Moreover, such $f(R)$ naturally provide an accelerated stage for the evolution of the universe, without to appeal to a dark energy component. ",
        "conclusions": "In this paper, we have studied the cosmological particle creation process to super-massive (nonrelativistic) real scalar particles in the framework of a $f(R)=R-\\beta/R^n$ modified gravity theory. The free parameters $n$ and $\\beta$ was constrained by Carvalho et al. 2008 \\cite{carvalho}, where the best-fit values for $n$ and $\\beta$ obtained from H(z)+BAO+CMB test are $0.03$ and $4.7$, respectively, with the parameters $n$ to lie in the intervals $n \\in [-0.25,0.35]$ and $\\beta \\in [2.3,7.1]$ at 99.7\\% c.l.. We have compared the results to the standard $\\Lambda CDM$ model, which corresponds to $n=0$ and $\\beta = 4.38$, and we have shown that, to super-massive particles, the effects of the modified gravity is to slightly increase the total particle number density as compared to the standard case at early times.  Such mechanism could be an alternative process to create dark matter particles in the universe. This opens the possibility that these particles could give rise to the dark matter in the early universe."
    },
    "1402/1402.2281_arXiv.txt": {
        "abstract": "The class of meteorites called carbonaceous chondrites are examples of material from the solar system which have been relatively unchanged from the time of their initial formation. These meteorites have been classified according to the temperatures and physical conditions of their parent planetesimals. We collate available data on amino acid abundance in these meteorites and plot the concentrations of different amino acids for each meteorite within various meteorite subclasses. We plot average concentrations for various amino acids across meteorites separated by subclass and petrologic type. We see a predominance in the abundance and variety of amino acids in CM2 and CR2 meteorites. The range in temperature corresponding to these subclasses indicates high degrees of aqueous alteration, suggesting aqueous synthesis of amino acids. Within the CM2 and CR2 subclasses, we identify trends in relative frequencies of amino acids to investigate how common amino acids are as a function of their chemical complexity. These two trends (total abundance and relative frequencies) can be used to constrain formation parameters of amino acids within planetesimals. Our organization of the data supports an onion shell model for the temperature structure of planetesimals. The least altered meteorites (type 3) and their amino acids originated near cooler surface regions. The most active amino acid synthesis likely took place at intermediate depths (type 2). The most altered materials (type 1) originated furthest toward parent body cores. This region is likely too hot to either favor amino acid synthesis or for amino acids to be retained after synthesis.  \\emph{Key Words:} astrobiology --- astrochemistry --- meteorites, meteors, meteoroids --- methods: data analysis --- molecular data ",
        "introduction": "Meteorites provide a geological and chemical record of the formation of organic molecules in the early solar system. Meteoritic parent planetesimals formed during the formation of the solar system and contain materials present at that time. Likely composed primarily of rock and ice, these parent bodies would be on the order of tens of kilometers in diameter. The closest analogues we have for these bodies are present-day asteroids. Meteorites found on Earth are likely fragments of their larger parent bodies, ejecta from collisions with parent bodies in the protoplanetary disk \\citep{Sephton2005,Gaidos2007}. Carbonaceous chondrites are thought to originate from type C asteroids, which are water-rich and likely came from the outer regions of the asteroid belt \\citep[and references therein]{Morbidelli2012}. There are a wide variety of meteorite types, all composed of a wide range of varying materials. Being inhomogeneous in nature, different compositions of meteorite are thought to indicate the degree of mixing of materials when their parent bodies were formed.  Amino acids that formed in meteoritic parent planetesimals are an important source of organic material that was delivered to young planets via meteoritic impacts. They are the basic building blocks in proteins used by life on Earth. This extraterrestrial origin of biomolecules suggests one possible way that the prebiotic conditions for life in the universe could have been established, both on our own early Earth and perhaps on other planets \\citep{Chyba1992,BottaBada,Sephton2002,Glavin2011}.  In this paper we investigate patterns in both total abundance and relative frequencies of amino acids in carbonaceous chondrites. Identification of these patterns offer insights into the physics behind their formation in planetesimals. Given the long time frame that parent bodies are likely to be in an aqueous state (a few million years \\citep{Travis2005}), it is possible to apply principles of equilibrium thermodynamics to follow the synthesis of amino acids. Total amino acid abundances depend on both temperature and the initial concentrations of reactants involved in the chemical processes that create them. Total abundance patterns therefore act as tracers of both temperature and molecular abundances within a parent body. In Paper II of this series (A. K. Cobb \\& R. E. Pudritz 2014, in preparation), we use thermodynamic principles and Gibbs free energies to model the synthesis of amino acids in a theoretical planetesimal.  Given some basic understanding of the formation of meteorites, it is now feasible to model meteoritic parent bodies in an effort to constrain their internal chemistry. \\citet{Young1999}, \\citet{Schulte2004} and \\citet{Travis2005} model parent bodies as agglomerations of ice, rock, and various interstellar organic material present during formation. The radionuclide decay of isotopes such as $^{26}$Al in their interiors keeps the temperature near the core high enough to melt water. Once liquid water is present, there are a variety of internal chemical processes that cause the organic material initially present to alter into new substances. These processes are collectively called aqueous alteration. See \\citet[and references therein]{Grimm1989,McSween1989,Weiss2013} for a review.  The class of meteorites called carbonaceous chondrites is of particular interest with regard to the origin of organic material. Representing approximately 4\\% of meteorite falls \\citep{Glavin2011}, these meteorites are relatively undifferentiated from the original material with which they originally formed. They contain high concentrations of water, carbon, and various organic materials. The carbonaceous chondrites are divided into subclasses and petrologic types based on their composition, mineralogy, and degrees of internal alteration. There is a range in alteration temperatures associated with different minerals found in meteorites, and these alteration temperatures are used to categorize meteorite samples. There are two types of alteration processes, aqueous and thermal. Carbonaceous chondrites that have undergone aqueous alteration show increased levels of organic material, including amino acids \\citep{Sephton2002,Weisberg2006,Glavin2011}.  Meteorite classification is based on two characteristics in particular: primarily, chemical composition, and secondarily, the amounts of aqueous and thermal processing. Aqueous alteration is a general term for the transformation of molecules to organic material in a hydrated environment - exactly what is occurring in the interior of meteoritic parent bodies resulting in the formation of amino acids. Thermal metamorphism occurs when alteration temperatures are too hot to retain a hydrated environment. See \\citet{Kojima1996}, \\citet{Weisberg2006} and \\citet{Weiss2013} for a review.  The range of temperatures that allows for significant amounts of this aqueous alteration is roughly $0^{\\circ}$C - $300 ^\\circ$C. At relatively high temperatures, $>300^\\circ$C, the parent bodies undergo thermal metamorphism instead; no aqueous alteration occurs, and we see little to no synthesis of organics. Petrologically speaking, more aqueously altered meteorites are given a classification value of type 1 or type 2, and thermally altered meteorites are listed as type 4 - 6. Type 3 meteorites are the most pristine samples of early solar system material, having experienced minimal degrees of either aqueous alteration or thermal metamorphism \\citep[and references therein]{Sephton2002,Weisberg2006}.  There is great diversity amongst carbonaceous chondrite literature regarding petrographic types and their corresponding alteration temperatures. These are well summarized in \\citet{Grimm1989}, \\citet{Morlok2013} and \\citet[and references therein]{Weiss2013}. See Section \\ref{classification} for a discussion of alteration temperature classification boundaries among petrographic types as well as the physical differences among meteorite classifications. Broadly speaking, the different petrographic types are classified based on compositional differences of meteorite samples. In each meteorite, minerals have undergone different degrees of melting or other alteration process, and each mineral has a unique melting point. Based on which minerals have been differentiated, a meteorite may be classified as having undergone alteration temperatures that reach the melting point of those specific minerals.   We explore the abundance patterns of $\\alpha$-proteinogenic amino acids found in carbonaceous chondrites. These are among the most common amino acids and are among the 20 used in our genetic code. We collate a large quantity of available data on amino acid concentration in a variety of carbonaceous chondrites from numerous sources and different laboratories. The data are shown graphically to more easily distinguish patterns in amino acid abundance across meteorite subclasses. In particular we use the data to identify patterns in amino acid concentration per meteorite class and per amino acid, as well as a total amino acid concentration per meteorite. We average total amino acid concentrations per meteorite class to observe overall abundance trends across the carbonaceous chondrites.  In addition to identifying abundance patterns across subclasses, we also investigate the relative frequencies of amino acids within two subclasses (CM and CR) of carbonaceous chondrites. \\citet{Higgs2009} examined a very limited set of carbonaceous chondrite data and found an empirical relation between the frequencies of some amino acids and their Gibbs free energies. The order they found related to the most common amino acids they called 'early-type', referring to the set of amino acids likely present on early Earth. These amino acids are glycine, alanine, aspartic acid, glutamic acid, valine, serine, isoleucine, leucine, proline, and threonine. Perhaps because of the availability of these amino acids due to their energetic favorability, they were incorporated into our genetic code. Increasing molecular complexity of amino acids is reflected in the increasing values of their associated Gibbs energies of formation. We expand on the amino acid landscape explored in \\citet{Higgs2009}.  Our treatment of the data is designed to emphasize patterns of amino acid synthesis in planetesimals across a range of alteration temperatures spanning meteorite subclasses. We also show the relative abundance patterns of amino acids in various meteorites based on environmental factors and internal chemistry during synthesis. Our analysis of these data shows that the CM and CR type 2 meteorites occupy a 'sweet spot' in alteration, which corresponds to specific meteorite classifications at which amino acid abundances peak. The patterns that emerge from the cumulative data plots reveal an optimal petrographic value of type 2, generally associated with temperatures between $0^{\\circ}$C and $100^{\\circ}$C. These values correspond to increased levels of aqueous alteration and the synthesis of amino acids. From these emerging patterns in abundance and frequency, we infer information about the chemical composition and temperatures at which amino acids would have likely been synthesized in parent planetesimals. The data collated here help provide insights into the chemistry behind the formation of these amino acids in parent bodies. ",
        "conclusions": "This work investigates amino acid abundance patterns and relative frequencies in $\\alpha$-proteinogenic amino acids of carbonaceous chondrites. We collated an extensive data set from the literature from a variety of sources covering a range of meteorites, including CI, CM, CR, CV, and CO, and petrologic types 1 - 3. Data have been organized so as to emphasize possible physical variables that control amino acid synthesis and abundance patterns. The data represent a wide variety of meteoritic parent body alteration states. Collation of this data allows for identification of patterns and the relationship between amino acid abundance and meteorite classification/petrologic type. We included amino acid abundances for 41 different samples of meteorite, covering five different subclasses of meteorite and three petrologic types.  It is possible to view our results on amino acid abundance patterns in the context of an onion shell model. We may view the sequence in types from 3 to 1 as descending from the cooler surface of a planetesimal, which has not been significantly aqueously altered, to a deeper layer, type 2, which has temperatures and conditions ideally suited to amino acid synthesis, from which the CM2- and CR2-types originate. With further descent toward the hot core, one finds conditions sampled by the type 1 chondrites, where temperatures are perhaps too hot to support amino acid presence. Detailed 3D simulations of planetesimals  by \\citet{Travis2005} support such a picture. Key results from our work are listed below.  \\begin{enumerate} \\item We present the data according to physical properties and trends in formation mechanisms. Type 1 meteorites, be they CI, CM, or CR, are the most aqueously altered specimens. Aqueous alteration during formation occurred at higher temperatures, possibly $50^{\\circ}$C to $150^{\\circ}$C, yet we do not observe a peak in total amino acid concentration at this location. All CI meteorites are petrologic type 1. There are only three listed type 1 CM/CR chondrites (two CM1 meteorites, one CR1 meteorite). \\item The type 2 meteorites (mostly CM-types and some CR-types) are generally associated with cooler alteration temperatures. These range broadly from $0^{\\circ}$C to $100^{\\circ}$C (based on splitting the CR temperature range into its petrographic types), show the greatest amino acid concentrations of all petrologic groups. The CM2 and CR2 chondrites contain the greatest quantities and variety of amino acids of all subclasses. \\item All CV and CO chondrites are type 3, associated with cool formation temperatures, roughly $0^{\\circ}$C to $50^{\\circ}$C. Type 3 meteorites are relatively unaltered from the time of their original formation. They have undergone minimal alteration, either aqueous or thermal metamorphism, and show a corresponding lack of significant amino acid presence. \\item These data indicate the existence of an optimal temperature range for the abundance of amino acids in planetesimals. The temperature range associated with type 2 chondrites, possibly ranging $0^{\\circ}$C to $100^{\\circ}$C, corresponds with amino acid abundances greater by several orders of magnitude. \\item For this optimal range, we identify the relative frequencies of amino acids relative to glycine for the CM2 and CR2 carbonaceous chondrites. In \\citet{Higgs2009}, relative frequencies are shown to occur according to thermodynamic principles - the more complex the molecule, the greater the Gibbs free energy necessary to synthesize it under certain conditions. The more extensive data shown here reinforces the conclusion that an underlying thermodynamic argument may be responsible. \\end{enumerate}  In A. K. Cobb \\& R. E. Pudritz (2014, in preparation), we investigate computationally if temperature and aqueous conditions in a planetesimal account for the overall abundance patterns we see in the data as well as for relative frequencies within subclasses. % \\newline  We would like to thank Daniel Glavin, Stephen Freeland, Boz Wing, Aaron Burton, Jamie Elsila, Jason Hein, and Darren Fernandes for stimulating discussions. We also thank an anonymous referee for several helpful suggestions that significantly improved the paper. A.K.C. would like to thank the Canadian Astrobiology Training Program who supported her with a CATP Graduate Fellowship, funded by NSERC CRSNG. R.E.P. was supported by an NSERC Discovery Grant.  \\newpage"
    },
    "1402/1402.4144_arXiv.txt": {
        "abstract": "Recent studies suggest that metal-poor stars enhanced in carbon but containing low levels of neutron-capture elements may have been among the first to incorporate the nucleosynthesis products of the first generation of stars. We have observed 16~stars with enhanced carbon or nitrogen using the MIKE Spectrograph on the Magellan Telescopes at Las Campanas Observatory and the Tull Spectrograph on the Smith Telescope at McDonald Observatory. We present radial velocities, stellar parameters, and detailed abundance patterns for these stars. Strontium, yttrium, zirconium, barium, europium, ytterbium, and other heavy elements are detected. In four stars, these heavy elements appear to have originated in some form of \\rpro\\ nucleosynthesis. In one star, a partial \\spro\\ origin is possible. The origin of the heavy elements in the rest of the sample cannot be determined unambiguously. The presence of elements heavier than the iron group offers further evidence that zero-metallicity rapidly-rotating massive stars and pair instability supernovae did not contribute substantial amounts of neutron-capture elements to the regions where the stars in our sample formed. If the carbon- or nitrogen-enhanced metal-poor stars with low levels of neutron-capture elements were enriched by products of zero-metallicity supernovae only, then the presence of these heavy elements indicates that at least one form of neutron-capture reaction operated in some of the first stars. ",
        "introduction": "\\label{introduction}   The first stars formed from only the products of Big Bang nucleosynthesis, yet the metals they produced and distributed into the interstellar medium forever changed the fundamental methods by which molecular gas clouds cool, collapse, and form stars. Cooling by fine structure line emission of ionized carbon (C~\\textsc{ii}), neutral oxygen (O~\\textsc{i}), and thermal emission from collisionally-excited dust grains is suspected to have enabled long-lived low-mass stars to form. Many of the most iron-poor stars known, with Fe/H ratios less than 10$^{-4}$ times the solar ratio, show C/H and O/H ratios of only 10$^{-3}$ to 10$^{-1}$. This may indicate that substantial enrichment in carbon or oxygen was closely linked with formation of the long-lived low-mass stars that are observable in the solar neighborhood. Theoretical studies of the nucleosynthesis reactions that may have occurred in the first stars have, understandably, focused their effort on production of metals from carbon through the iron group. Elements heavier than the iron group, here considered to be those with $Z >$~32, are difficult to detect observationally in the most iron-poor stars and almost certainly had no effect on subsequent star formation due to their low abundances of 10$^{-13}$ or less per hydrogen atom.  Carbon-enhanced metal-poor stars with no enhancement of neutron-capture elements are commonly referred to as members of the ``CEMP-no'' class, using the classification scheme proposed by \\citet{ryan05} and \\citet{beers05}. \\citeauthor{ryan05}\\ proposed that these stars formed from gas clouds pre-enriched with high levels of carbon by previous generations of supernovae, rather than acquiring the carbon enrichment by mass transfer from an evolved companion. The CEMP-no stars include three of the four most iron-poor stars known \\citep{christlieb02,christlieb04,bessell04, frebel05,frebel06,frebel08,aoki06,norris07,norris12}. CEMP-no stars are frequently enhanced in nitrogen (N, $Z =$~7), oxygen (O, $Z =$~8), sodium (Na, $Z =$~11), magnesium (Mg, $Z =$~12), aluminum (Al, $Z =$~13), and silicon (Si, $Z =$~14). A few CEMP-no stars are found in binary systems, but the binary nature of the ensemble of CEMP-no stars is clearly unlike that of the CEMP stars enhanced in material produced by the slow neutron-capture process (the \\spro) in a companion that passed through the thermally-pulsing asymptotic giant branch (TP-AGB) phase of evolution.  \\citet{norris13} summarize these properties and weigh the evidence that stars in the CEMP-no class were among the first stars to incorporate metals produced by zero-metallicity stars. This evidence includes the chemically-primitive stars in ultra-faint dwarf galaxies (e.g., \\citealt{frebel10,norris10a,norris10b,simon10,koch13}), including the discovery of a CEMP-no star in each of Segue~1 \\citep{norris10a} and Bo\\\"{o}tes~I \\citep{lai11,gilmore13}; the greater chemical inhomogeneity, on average, of field stars that are kinematically associated with the outer halo (e.g., \\citealt{fulbright02,stephens02,gratton03, roederer09a,ishigaki10,ishigaki13,nissen10}); the increasing fraction of carbon-enhanced stars with decreasing metallicity and increasing distance above the Galactic plane \\citep{frebel06,carollo12}; and the low frequency of carbon-enhanced metal-poor damped Lyman-$\\alpha$ (DLA) systems at redshifts 2~$< z <$~6.3, hinting that an epoch dominated by carbon-enhanced systems, if one existed at all, must have occurred at even higher redshifts (e.g., \\citealt{cooke11a,cooke11b,becker12}; but note that \\citeauthor{becker12}\\ express reservations about the high level of carbon enhancement in the DLA reported by \\citeauthor{cooke11a}). These links between carbon enhancement, low metallicity, and remote environments hint that carbon-enhanced stars formed in chemically-primitive regions that experienced relatively few enrichment events.  What kind of stars were responsible for prodigious carbon production in the early Universe? The abundance patterns found in the CEMP-no stars are consistent with model predictions for zero-metallicity stars that were massive and rotating \\citep{fryer01,meynet06,meynet10}, underwent faint ``mixing and fallback'' supernova explosions \\citep{umeda03,umeda05,tominaga13}, or whose supernovae had relativistic jets \\citep{tominaga07}. These zero-metallicity supernovae are predicted to have seeded pristine gas clouds with carbon, oxygen, and other metals, enabling low-mass star formation to occur.  We now pose a related question:\\ \\textit{What heavy element ($Z >$~32) abundance patterns are found in stars in the CEMP-no class?} We address this question using seven metal-poor stars identified in the abundance survey of \\citet{roederer14}. These stars are carbon-enhanced and barium-poor (Ba, $Z =$~56). We supplement these stars with nine nitrogen-enhanced metal-poor stars that are barium-poor (defined here as NEMP-no, cf.\\ \\citealt{johnson07}) from the same survey. Our goal is not to resolve whether the CEMP-no or NEMP-no stars are the immediate descendants of zero-metallicity stars. Instead, we aim to characterize the heavy element abundance signatures in these stars to motivate studies of neutron-capture nucleosynthesis in stars with zero or extremely low levels of metals. Unfortunately, existing nucleosynthesis calculations for such stars rarely extend to nuclei heavier than the iron group. For example, the reaction network of \\citet{tominaga13} is truncated at bromine (Br, $Z =$~35), which is several mass units short of the lightest neutron-capture element commonly studied in metal-poor stars, strontium (Sr, $Z =$~38).  Throughout this work we adopt the standard definitions of elemental abundances and ratios. For element X, the logarithmic abundance is defined as the number of atoms of X per 10$^{12}$ hydrogen atoms, $\\log\\epsilon$(X)~$\\equiv \\log_{10}(N_{\\rm X}/N_{\\rm H}) +$~12.0. For elements X and Y, the logarithmic abundance ratio relative to the solar ratio, denoted [X/Y], is defined as $\\log_{10} (N_{\\rm X}/N_{\\rm Y}) - \\log_{10} (N_{\\rm X}/N_{\\rm Y})_{\\odot}$. Abundances or ratios denoted with the ionization state indicate the total elemental abundance as derived from transitions of that particular state after ionization corrections have been applied. When reporting relative abundance ratios for elements X and Y, these ratios compare the total abundances of X and Y derived from like ionization states; i.e., neutrals with neutrals and ions with ions. ",
        "conclusions": ""
    },
    "1402/1402.4833_arXiv.txt": {
        "abstract": "We study the feasibility of observationally constraining the rotation rate of hot Jupiters, planets that are typically assumed to have been tidally locked into synchronous rotation.  We use a three-dimensional General Circulation Model to solve for the atmospheric structure of two hot Jupiters (\\HDo~and \\HDt), assuming rotation periods that are 0.5, 1, or 2 times their orbital periods (2.2 and 3.3 days, respectively), including the effect of variable stellar heating.  We compare two observable properties: 1) the spatial variation of flux emitted by the planet, measurable in orbital phase curves, and 2) the net Doppler shift in transmission spectra of the atmosphere, which is tantalizingly close to being measurable in high-resolution transit spectra.  Although we find little difference between the observable properties of the synchronous and non-synchronous models of \\HDo, we see significant differences when we compare the models of \\HDt.  In particular, the slowly rotating model of \\HDt~has an atmospheric circulation pattern characterized by westward flow and an orbital phase curve that peaks after secondary eclipse (in contrast to all of our other models), while the quickly rotating model has a net Doppler shift that is more strongly blue-shifted than the other models.  Our results demonstrate that the combined use of these two techniques may be a fruitful way to constrain the rotation rate of some planets, and motivate future work on this topic. ",
        "introduction": "Perhaps the most commonly used assumption in studying hot Jupiters is this: that the planets with zero eccentricity have been tidally locked into synchronous rotation, meaning that their orbital and rotational periods are equal, with one hemisphere in unending day and the other in perpetual night.  This assumption is based on timescale arguments that tidal forces should lock a planet into a state of synchronous rotation long before they circularize its orbit \\citep[e.g.,][]{Rasio1996}.  However, the physical mechanism of tidal dissipation within a gas giant is a complex process and an area of active research \\citep[e.g.,][]{Ogilvie2004,Remus2012,Weinberg2012}.  Recently \\citet{Socrates2012} have argued that the simplifying assumptions commonly used to model the tidal response of hot Jupiters may not accurately capture the physics at work, especially when it comes to the tidal circularization of highly eccentric gas giants.  Additionally, even if a planet were tidally locked into a synchronous state, the flow of angular momentum between the planet's orbit, atmosphere, and/or interior could torque the planet out of synchronous rotation \\citep{Showman2002}.  In particular, \\citet{Arras2010} have proposed that the ``thermal tide'' torque originally used to explain Venus' spin state \\citep{Gold1969} could be at work in hot Jupiters, resulting in a steady state with asynchronous rotation and a source of internal heating through the constant dissipation of gravitational tides.  Since many hot Jupiter radii are observed to be larger than can be explained without an extra source of heating inside the planet,\\footnote{\\citet{Spiegel2013} contains a good, recent description of this open question and compares the proposed solutions.} it is especially important to question whether there are any observable signatures that could be used to identify non-synchronous rotation of hot Jupiters.  A gas giant planet locked into synchronous rotation can still support differential fluid motions throughout its atmosphere and interior; the synchronicity is a requirement on the total angular momentum of the body.  If a solid core exists within a synchronously rotating planet, it will presumably be locked into solid-body rotation at the synchronous rate and dominate the (non-orbital) angular momentum of the planet.  The core is in continuous contact with the fluid layers above and angular momentum transport will ensure that motion in the upper atmospheric layers is influenced by the global rotation of the planet.  The atmospheric patterns of winds and (observable) temperature structures are strongly dependent on the rotation rate, through the influence of the Coriolis force on atmospheric dynamics.  Given the assumption of synchronous rotation, the Coriolis force should have a much weaker effect on the atmospheres of hot Jupiters ($\\Prot \\sim 3$ days) than on the atmosphere of the more quickly rotating namesake, Jupiter ($\\Prot \\sim 10$ hours), and early estimates of dynamical quantities predicted that this should result in very large atmospheric features on hot Jupiters \\citep{Showman2002,Cho2003,Menou2003}.  This has been born out in three-dimensional General Circulation Models (GCMs) of the atmospheres of synchronously rotating hot Jupiters \\citep{Showman2009,Thrastarson2011,Perna2012,Rauscher2012b,DobbsDixon2013,Mayne2014}.  One common atmospheric characteristic found in most GCMs is a super-rotating\\footnote{``Super-rotating'' means that the flow is eastward, with greater angular momentum than the planet's rotation at that radius and latitude.} equatorial jet, which results in the hottest region of the atmosphere being advected to the east of the substellar point.  The development of the eastward jet is a general property of the hemispheric forcing regime of hot Jupiters \\citep{Showman2011}, although it can be inhibited by the presence of high viscosity or strong drag \\citep{DobbsDixon2010,Perna2010a,Showman2011,Rauscher2013}, in which case the hottest region remains near the substellar point.  The lack of a shift from the substellar point does not necessarily preclude the existence of an equatorial jet, however, since in the case of very short radiative timescales (i.e., for very hot atmospheres), the gas at the substellar point cools more quickly than winds can advect it away from the substellar point \\citep{Perna2012,Showman2013}.  For a few of the brightest hot Jupiters it has been possible to measure the flux emitted from the planet as a function of longitude, by observing the planetary system throughout an orbit and watching for the varying flux as different phases come into view.  Of the measured orbital phase curves for hot Jupiters on circular orbits, we see the expected range of behavior:\\footnote{Note that the westward phase shift seen in the orbital phase curve for the hot Jupiter Kepler-7b was observed at optical wavelengths and is interpreted as reflected light from spatially inhomogeneous clouds \\citep{Demory2013}, which is not inconsistent with the presence of an eastward equatorial jet.} some planets have their brightest regions well aligned with the substellar point \\citep{Cowan2007,Snellen2009,Welsh2010}, while others have the brightest region shifted to the east of the substellar point \\citep{Cowan2012,Knutson2012}, including the curious case of $\\upsilon$~Andromedae~b \\citep{Crossfield2010}, which has a surprisingly large shift ($\\sim$80\\degrees).\\footnote{Naively we would expect the amplitude of the flux variation to be anti-correlated with the amount by which the hot spot is shifted, since the shift is a result of winds transporting the hot gas faster than it can cool, but such efficient winds should result in a more homogenous horizontal temperature structure \\citep[e.g.][]{Cowan2011a}.  No current GCM can reproduce the $\\upsilon$ And b phase curve.}  The planet with the best signal-to-noise observations, \\HDo, also has eclipse mapping measurements that independently support the eastward shift of the brightest region \\citep{Majeau2012}, although see \\citet{deWit2012} for a careful discussion of parameter degeneracies.  \\citet{Showman2009} performed the first simulations of a non-synchronously rotating hot Jupiter, modeling the planet \\HDo.  Although there was some variation in the circulation pattern for different rotation rates---which we will return to discuss in detail below---all models developed an eastward jet at the equator and produced orbital phase curves in which the brightest region was shifted east of the substellar point, although the magnitude of the shift varied with the rotation rate \\citep[at a level comparable to the shift in the peak emission between models run at different metallicities,][]{Showman2009}.  Their results imply that an orbital phase curve measurement cannot conclusively differentiate between synchronous and non-synchronous rotation.  Finding ways to measure the rotation rate of solar system giant planets has been a challenge in its own right.  The most widely accepted values come from \\emph{in situ} satellite measurements of the periodicity of the planetary magnetic field (assumed to be generated in the deep interior and therefore tied to that bulk rotation rate) or modulation of radio emission, presumed to be tied to the rotating magnetic field lines \\citep[e.g.,][]{Desch1981,Desch1986,Warwick1989}.  However, additional types of measurements have also been used to estimate the rotation rates, such as: the oblateness of a planet, Doppler-broadening of line widths in atmospheric spectra, periodicity in light curves, and tracking of imaged cloud features \\citep[see a discussion in the review by][]{Stevenson1982}.  The rotation and global structure of the solar system giant planets continues to be an area of active research; one example being the recent work by \\citet{Helled2010}, which calls into question whether the magnetic rotation periods are tied to the bulk rotation rates of Uranus and Neptune, since those periods do not match the ones derived from fitting to the planets' oblateness measurements.  It may be a long wait before we can achieve \\emph{in situ} measurements of exoplanet properties, but the other methods of measuring planetary rotation rates may be possible in the nearer future.\\footnote{The first measurement of an exoplanet rotation rate was published during the refereeing process for this paper.  \\citet{Snellen2014} observed rotational broadening in their high-resolution spectrum of the directly imaged young exoplanet $\\beta$~Pictoris~b, corresponding to an equatorial rotational velocity of 25$\\pm$3 \\kms.}  There is a long history of searching for radio emission from exoplanets, which has so far resulted only in upper limits \\citep[aside from a very recent possible first detection,][]{Lecavelier2013}.  In order to use radio emission to estimate rotation rates a modulated signal is needed, which may only be detectable with current instruments for strong planetary magnetic field strengths and fortuitous viewing geometry \\citep{Hallinan2013}.  Even if a periodic signal were detected, it might be difficult to confidently identify the period as the rotation period of the planet, rather than some stellar period or a period related to star-planet magnetic interaction.  Rotational periods could potentially be measured for Earth-like exoplanets by observing variation in the reflected light from the planet with next-generation instruments \\citep{Palle2008}, but for close-in planets the periodic variation in thermal emission is subject to several parameters, including the radiative timescale of the atmosphere and wind speeds on the planet \\citep{Cowan2011a}.  For hot Jupiters we expect the period of variation to be well matched to the orbital period, regardless of the rotation rate \\citep[based on the work of][discussed above]{Showman2009}.  The oblateness of a planet could possibly be measured by the detailed shape of the light curve as the planet transits its star \\citep{Seager2002,Barnes2003}; however, a recent observational limit could only constrain the rotation period of the hot Jupiter HD 189733b to be greater than about 10 hours, or $\\sim$20\\% of the synchronous rotation period \\citep{Carter2010}.  Here we investigate the method of using Doppler shifts in a hot Jupiter's transit spectrum to observationally constrain its rotation rate.  Previous work on this topic has shown that the Doppler rotational broadening in the spectrum may be of comparable magnitude to the expected shifts from atmospheric winds \\citep{Brown2001,Spiegel2007,Kempton2012,Showman2013}, meaning that both effects must be included in any predictive model.  In this paper we compare models for hot Jupiters at different rotation rates, consistently including the atmospheric circulation structure calculated from a 3D GCM using that rotation rate.  In Section~\\ref{sec:model} we discuss our modeling framework, both the radiative transfer with Doppler shifts \\citep[from][]{Kempton2012} and the GCM \\citep[from][]{Rauscher2012b}, now adapted to allow for non-synchronous rotation.  In Section~\\ref{sec:circ} we present the results of the GCMs and discuss the variation in circulation patterns between models.  We show the observable properties of each model in Section~\\ref{sec:obs}: the thermal phase curves and the transit spectra.  Finally, in Section~\\ref{sec:conc} we summarize our results and discuss the potential of observationally constraining hot Jupiter rotation rates. ",
        "conclusions": "\\label{sec:conc}  We have presented a study of the observational characteristics from models of hot Jupiters that are not in synchronous rotation states.  We compared models with parameters chosen to represent two well known hot Jupiters, \\HDo~and \\HDt, the former having a slightly faster synchronous rotation rate than the latter.  For each planet we ran consistent 3D circulation models with rotation rates 0.5, 1, and 2 times its orbital rate ($\\Porb=\\Prot$ is the standard assumption of a tidally locked, synchronous state).  We found the standard development of an eastward equatorial jet in all models but the slowly rotating model of \\HDt, in which the circulation was instead primarily westward.  This indicates the presence of a circulation regime transition at very low rotation rate, although a more detailed characterization and understanding of the regime shift is warranted as a topic for future study.  The observational characteristics we computed from each model were: 1) the thermal emission that would be observed in an orbital phase curve and 2) the net Doppler shift that would be recovered from high-resolution spectra taken during transit.  Neither type of measurement showed much potential as a means by which to constrain the rotation rate of \\HDo, as all three models produced very similar signatures.  For \\HDt, however, we found that the slowly rotating model would be distinguishable by a delayed peak in the orbital phase curve, while the quickly rotating model would uniquely present a more strongly blue-shifted Doppler signal.  Our results show that in some cases these observational methods could be used to constrain the rotation rate of hot Jupiters, but that further work will be required to determine in more detail the system parameters that enable the use of these techniques.  By modeling a broader swath of parameter space, we can hope to more clearly identify the reasons that our models of \\HDt~show observable differences with rotation rate, while our models of \\HDo~do not.  We also caution that there are several modeling particularities that could influence the precise values of the Doppler-shift or the phase of peak flux that our models predict.  For example, the strength of numerical dissipation that is applied to a model, necessary to prevent the build-up of small scale noise, is directly related to the wind speeds calculated in the model \\citep{Heng2011b} and there are no calculations based on physical conditions that can prescribe what value should be used \\citep[see][for some of the complexity involved]{Thrastarson2011}.  Similarly, the temperature structure of the planet and the shape of the orbital phase curve can depend on the particular constituents of the atmosphere.  Specifically, initial data for the planet \\HDt~indicated that the atmosphere of this planet may be enhanced in an optical absorber, leading to a temperature inversion in its upper atmosphere \\citep[e.g.,][]{Burrows2007b}.  However, recent additional data potentially weaken the evidence for a temperature inversion \\citep{Zellem2014}.  Models with stratospheric absorbers have different phase curves and sometimes have different transit spectra properties \\citep{Fortney2006,Burrows2010,Fortney2010}, but we do not include any extra absorbers in our models of \\HDt.  Finally, in these models we have neglected the magnetic effects that result from the finite thermal ionization in the hot atmospheres and the presence of a planetary magnetic field \\citep{Perna2010a,Perna2010b}.  By doing so, we may be missing an important piece of physics, because we have previously shown that these effects can influence atmospheric circulation, particularly for \\HDt~(which is hotter and therefore more ionized than \\HDo), and alter the observed thermal phase curves \\citep{Rauscher2013} and net Doppler shifts \\citep{Kempton2012}.  However, the inclusion of magnetic effects in atmospheric circulation models is still in its infancy, with the MHD model by \\citet{Batygin2013} indicating that non-zero magnetic fields should preclude day-to-night flow (that all winds should be in the east-west direction), while \\citet{Rogers2014} found that the circulation in their (cooler) MHD atmospheres retained the familiar patterns seen in standard GCMs, albeit with slower wind speeds.  The rotation rate of a planet is one of its fundamental characteristics and yet can be very difficult to measure.  We have shown that in some cases in may be possible to combine multiple observational techniques in order to constrain the rotation rate of hot Jupiters.  It would be particularly valuable to measure rotation for this class of exoplanet, since they are typically assumed to have been tidally locked into synchronous rotation and an observational constraint could help to inform theories of tidal dissipation in these systems.  Although more work will be necessary before a set of observations could be uniquely translated into a measured rotation rate, we have demonstrated that this is a fruitful line of theoretical investigation and that the development of more sensitive observations is well motivated."
    },
    "1402/1402.6231_arXiv.txt": {
        "abstract": "We investigate how a stochastic gravitational wave background, produced from a discrete set of astrophysical sources, differs from an idealised model consisting of an isotropic, unpolarised, and Gaussian background. We focus, in particular, on the different signatures produced from these two cases, as observed in a cross-correlation search. We show that averaged over many realisations of an astrophysical background, the cross-correlation measurement of an astrophysical background is identical to that of an idealised background. However, any one realisation of an astrophysical background can produce a different signature. Using a model consisting of an ensemble of binary neutron star coalescences, we quantify the typical difference between the signal from individual realisations of the astrophysical background and the idealised case. For advanced detectors, we find that, using a cross-correlation analysis, astrophysical backgrounds from many discrete sources are probably indistinguishable from an idealised background. ",
        "introduction": "One of the science goals of second-generation gravitational-wave detectors such as Advanced LIGO~\\cite{aligo} and Virgo~\\cite{virgo} is to detect a stochastic gravitational-wave background. A stochastic background arises from the superposition of many gravitational-wave sources, each of which cannot be individually resolved~\\cite{Allen97,2000PhR...331..283M}. A stochastic background can be created in the early universe following inflation~\\cite{grishchuk,starob,eastherlim,peloso}, during a phase transition~\\cite{2000PhR...331..283M}, or from cosmic strings~\\cite{caldwellallen,DV1,DV2,cosmstrpaper,olmez1,olmez2} to name a few scenarios. Less speculative astrophysical stochastic backgrounds are expected to arise from more vanilla objects such as compact binaries~\\cite{phinney,kosenko,zhu,StochCBC,RegPacCBC}, neutron stars~\\cite{cutler,RegMan,rosado_ns,howell2010,marassi,RegPac,owen,barmodes2,barmodes3}, core collapse supernovae~\\cite{firststars,howell2004,buonanno2005,marassi2009}, white dwarf binaries~\\cite{phinney_whitedwarfs} and super-massive black hole binaries~\\cite{wyithe,jaffe,enoki}.  A stochastic background can be described in terms of its energy density spectrum $\\Omega_\\text{gw}(f)$, which is the fractional contribution of the energy density in gravitational waves relative to the total energy density needed to close the universe~\\cite{Allen97}:  \\begin{equation} \\Omega_\\text{gw}(f) = \\frac{1}{\\rho_c}\\frac{d\\rho_\\text{gw}}{d\\ln f} . \\end{equation}  \\noindent Here $\\rho_c$ is the critical energy density of the universe and $d \\rho_\\text{gw}$ is the gravitational-wave energy density between $f$ and $f+df$. Typically, searches for a stochastic background estimate $\\Omega_\\text{gw}(f)$ using a cross-correlation statistic (see, e.g.,~\\cite{Allen97,stoch-S5}), which we denote $\\widehat{Y}(f)$. In~\\cite{Allen97}, the estimator $\\widehat{Y}(f)$ is derived for the case of an isotropic, unpolarised, and Gaussian background. While subsequent work has relaxed the assumption of isotropy~\\cite{PhysRevD.80.122002,PhysRevLett.107.271102}, it is still typically assumed that the observed background is Gaussian, (see, e.g.,~\\cite{stoch-S5}). However, it is likely that the first detection of a stochastic background will be of a non-Gaussian background of astrophysical origin~\\cite{StochCBC}. Non-Gaussian backgrounds exhibit fluctuations arising from the discrete nature of their composition; no two realisations are exactly the same.  In this paper we investigate how the non-Gaussianity of astrophysical stochastic backgrounds affects cross-correlation measurements of $\\Omega_\\text{gw}(f)$. First, we calculate $\\langle \\widehat{Y}(f)\\rangle_{n,h}$, the expectation value of $\\widehat{Y}(f)$ in the presence of a non-Gaussian background averaged over both realisations of detector noise and realisations of an astrophysical stochastic background. The answer, we show, is identical to the case of an isotropic Gaussian background. Next, we calculate $\\langle \\widehat{Y}(f)\\rangle_n$, the expectation value of $\\widehat{Y}(f)$ averaged over realisations of detector noise but considering only {\\em a single realisations of an astrophysical background}. The answer, this time, is different than the case of an isotropic Gaussian background. By comparing these two calculations, we characterise the signature caused by the discreteness of astrophysical backgrounds. We proceed to estimate the size of this signature in upcoming observations by advanced detectors.  The remainder of the paper is organised as follows. In section~\\ref{formalism}, we review the procedure for a cross-correlation search for a stochastic background (subsection~\\ref{review}), characterise the statistical behaviour of astrophysical backgrounds (subsections~\\ref{ensemble} and~\\ref{realisation}), and introduce a novel formalism for characterising astrophysical backgrounds. Then, in section~\\ref{numerical}, we present the results of a numerical investigation, which quantifies the statistical fluctuations between different realisations of the stochastic background. Finally, in section~\\ref{conclusions}, we summarise our results and discuss the implication for future gravitational-wave observations. ",
        "conclusions": "Many previous studies of the stochastic gravitational-wave background have assumed a signal that is isotropic, unpolarised, and Gaussian. However, non-Gaussian backgrounds from compact binary coalescence represent one of the most exciting sources for second-generation detectors such as Advanced LIGO and Virgo. In this paper, we investigated the statistical properties of stochastic backgrounds originating from a discrete set of astrophysical sources and how they will appear in future cross-correlation searches. In the course of our investigation, we found it useful to define a novel description of astrophysical backgrounds: a discrete overlap reduction function. We find that the discreteness of astrophysical backgrounds is unlikely to produce a measurable bias in upcoming observations by second-generation detectors.  Here we focused on upcoming advanced detectors observing a population of binary neutron star sources. However, we note that the situation may be more complicated for the proposed third-generation Einstein Telescope \\cite{ETMDC1}. In particular, we raise the possibility that the removal of above-threshold binary events may create a selection bias. This is because we expect face-on events, directly above the detector, will be preferentially detected compared to events with less favourable orientations and locations, which in turn, may create an apparent anisotropy. The effect may be more pronounced for the Einstein telescope (with only one detector) versus a network of 2--5 advanced detectors. Future work will characterise the magnitude of this effect for the Einstein Telescope."
    },
    "1402/1402.3308_arXiv.txt": {
        "abstract": "CO is thought to be a vital building block for prebiotic molecules that are necessary for life. Thus, understanding where CO existed in a solid phase within the solar nebula is important for understanding the origin of life. We model the evolution of the CO snow line in a protoplanetary disk. We find that the current observed location of the CO snow line in our solar system, and in the solar system analogue TW Hydra, cannot be explained by a fully turbulent disk model. With time-dependent disk models we find that the inclusion of a dead zone (a region of low turbulence) can resolve this problem. Furthermore, we obtain a fully analytic solution for the CO snow line radius for late disk evolutionary times. This will be useful for future observational attempts to characterize the demographics and predict the composition and habitability of exoplanets. ",
        "introduction": "The most abundant volatiles in a protoplanetary disk are $\\rm CO$, $\\rm CO_2$ and $\\rm H_2O$. A snow line marks a radial location in a disk where the mid-plane temperature\\footnote{We assume that the gas and dust temperatures in the disk are equal. This assumption may break down in the upper layers of the disk, but is reasonable at the disk midplane \\citep[e.g.][]{Dullemond2007}.} drops sufficiently so that a volatile condenses out of the gas phase to become solid\\footnote{Volatiles condense on to dust grains to become solid \\citep{Tielens1982} rather than forming pure ice.}. Thus, each volatile has a different snow line radius, water ice being closest to the host star, then $\\rm CO_2$ and then $\\rm CO$. Snow lines are thought to regulate the planet formation process \\citep[e.g.][]{Oberg2011}. Giant planets, for instance, are expected to form outside the water snow line because the density of solids is significantly higher outside of this radius \\citep[e.g.][]{Pollack1996,Morales2011,Ros2013}. Snow lines are also important because dust piles up in the pressure trap just inside the snow line and the grains become stickier as ice condenses on their surfaces. The composition of a planet and its atmosphere are determined by where the planet forms, and where it accretes material relative to the snow lines \\citep[e.g.][]{Oberg2011}.  The water snow line occurs at a temperature of around $170\\,\\rm K$ \\citep{Hayashi1981} and currently in our solar system is observed to be at a radius of $2.7\\,\\rm AU$, within the asteroid belt \\citep{Abeetal2000,Morbidellietal2000,MartinandLivio2013asteroids}. The CO snow line occurs at a cooler temperature of about $T_{\\rm CO,snow}=17\\,\\rm K$ \\citep{Oberg2005}.  Comets, from the Kuiper belt, have varying amounts of CO, suggesting that they formed close to the CO snow line \\citep{Ahearnetal2012}.  Some dwarf planets such as pluto and comets in the Kuiper belt contain the even more volatile $\\rm N_2$ gas \\citep[e.g.][]{Cochran2000} implying that they may have formed beyond the CO snow line. The Kuiper belt is currently thought to have formed in the approximate region $27 -35\\,\\rm AU$ \\citep{Levison2008} and thus, at the time of planetesimal formation the CO snow line would have been in this region. The CO snow line could mark the transition from the planet forming region to the formation of smaller icy bodies and dwarf planets, like Pluto.  The water snow line is hard to observe in exosolar systems because it is very close to the star. However, the CO snow line is an easier target because it is farther away. The best studied CO snow line is in a solar nebula analogue disk around TW Hya. The star has a mass of $0.8\\,\\rm M_\\odot$ and an age of less than $10\\,\\rm Myr$ \\citep{Hoffetal1998}. \\cite{Qietal2013} observed the reactive ion $\\rm N_2H^+$ which is only present where CO is frozen out. They found that the CO snow line lies at $28 -31\\,\\rm AU$, similar to that in our solar system.  Understanding the evolution of the CO snow line is essential to deciphering the origins of prebiotic molecules that are necessary for life \\citep{Tielens1997}. CO ice is needed to form methanol which is a building block for more complex organic molecules. Comets are thought to have bombarded the early Earth, thus delivering these molecules to Earth and allowing life to emerge. In this work we therefore consider the evolution of the CO snow line in various models of protoplanetary disks.  Angular momentum transport in protoplanetary disks is thought to be driven by turbulence generated by the magneto-rotational instability (MRI) \\citep{BH1991}.  However, it is now widely acknowledged that protoplanetary disks are not sufficiently ionised for the MRI to operate throughout. They contain a {\\it dead zone}, a region of low turbulence at the disk mid-plane where the MRI is suppressed \\citep[e.g.][]{Gammie1996}.  In the present Letter, we investigate the evolution of the CO snow line in disks with and without a dead zone. ",
        "conclusions": "\\label{discussion}   Shearing box simulations suggest that a small amount of turbulence may be driven in the dead zone by the turbulence in the disk surface layers \\citep[e.g.][]{Fleming2003, Simonetal2011}.  The addition of this small viscosity is unlikely to suppress outbursts \\citep{MartinandLubow2013dza}, although the triggering may be due to heating from the additional turbulence, rather than self gravity \\citep{Baeetal2013}.  In the limit of small active layer surface density, the steady state dead zone solution would be the same as the fully MRI turbulent solution in Section~\\ref{turb} but with a smaller $\\alpha_{\\rm m}$. In the figures we have also considered a smaller turbulence of $\\alpha_{\\rm m}=10^{-4}$ for this comparison. We find that, with an $\\alpha_{\\rm m}$ that is two orders of magnitude smaller than that in the active layers, it remains difficult to explain the current location of the CO snow line in our solar system, and in TW Hya.  The CO snow line in the disk around the Herbig~Ae star HD 163296 has been found to lie at a radius of around $155\\,\\rm AU$ \\citep{Qietal2011}. With observed parameters of $M_\\star=2.3\\,\\rm M_\\odot$, $T_\\star=9333\\,\\rm K$, $R_\\star=2\\,\\rm R_\\odot$ and $\\dot M=7.6\\times 10^{-8}\\,\\rm M_\\odot\\,yr^{-1}$, our fully turbulent disk model (equation~\\ref{rturb}) predicts a CO snow line radius of $37\\,\\rm AU$. The dead zone model (equation~\\ref{app}) predicts a radius of $62\\,\\rm AU$. Thus, neither model can explain such a large CO snow line radius (although the model with a dead zone shifts the radius in the right direction). We suggest that disk flaring could account for this. Approximations for the temperatures of flared disks would predict a CO snow line radius $>100\\,\\rm AU$ \\citep[see Fig. 4 in][]{Chiang1997}. In contrast, the inner regions of the disk in TW~Hya are flat, although it is moderately flared at radii $R>45\\,\\rm AU$. Thus, flaring does not affect the CO snow line radius in TW~Hya. Furthermore, this would suggest that our solar nebula was not flared, at least in the inner regions.    We have found that a fully MRI turbulent disk predicts a CO snow line that is much closer to the host star than that observed in our solar system and in the solar neubula analogue TW Hya. With a dead zone, a small amount of self-gravity heats the more massive disk and the CO snow line radius is moved outwards (in agreement with the observations). We have also found a fully analytic solution for the snow line radius in a disk with a dead zone for low infall accretion rates, appropriate for the later stages of protoplanetary disk evolution. The solution is valid providing that the surface density ionised by external sources, $\\Sigma_{\\rm crit} \\ll 100\\,\\rm g\\,cm^{-2}$. This formula could prove useful for determining composition and habitability of exo-solar planets."
    },
    "1402/1402.5624_arXiv.txt": {
        "abstract": "{% The aim of our contribution is to shed some light on open questions facing the high density nuclear many-body problem. We focus our attention on the conceptual issue of naturalness and its role for the baryon-meson coupling for nuclear matter at high densities. As a guideline for the strengths of the various couplings the concept of naturalness has been adopted. In order to encourage possible new directions of research, we discuss relevant aspects of a relativistic effective theory for nuclear matter with {\\it natural} parametric couplings and genuine many-body forces. Among other topics, we discuss in this work the connection of this theory with other known effective Quantum Hadrodynamics (QHD) models found in  literature and how we can potentially  use our approach to  describe new physics for neutron stars. We also show some preliminary results for the equation of state, population profiles and mass-radius relation for neutron stars assuming local charge neutrality and beta equilibrium. } ",
        "introduction": "Quantum Chromodynamics (QCD) with quark and gluon degrees of freedom is the underlying theory of the strong interaction physics. However, at energies relevant for most nuclear phenomena, hadron and meson fields provide a more appropriate and convenient description. In this domain, effective quantum field theory (EFT) and, more precisely, Quantum Hadrodynamics (QHD) represent efficient frameworks for describing the nuclear many-body problem as a relativistic system of baryons and mesons.  The motivation for using QHD in this domain can be synthesized in a few words: simplicity, efficiency, consistency. There are many arguments that vindicate this assertion: it is the simplest approach for describing nuclear matter and nuclear systems based on a local Lorentz-invariant lagrangian density; it represents an efficient way to parameterize an S matrix or other observables in nuclear matter; this approach is consistent with analyticity, unitarity, causality, cluster decomposition, parity conservation, spin and isospin symmetries, and spontaneously breaking of chiral symmetry for systems undergoing strong interaction.  Moreover, since the most relevant global phenomena in nuclear physics are confined to longer length scales, ie, in the interaction sector where less massive meson fields do\\-mi\\-na\\-te, it is not necessary to explicitly include dynamics at significantly shorter length scales, where the presence of the most massive meson fields becomes crucial.  Thus, in describing global properties of nuclear systems through effec\\-ti\\-ve field theory, one can ignore the latter by {\\it in\\-te\\-grating out} heavier degrees of freedom corresponding to shorter length scales. It is important to notice, however, the effects of these heavier degrees of freedom are still implicitly contained in the coupling parameters of the effective field theory (Serot \\& Walecka ~1997). EFT and QHD contain on the other hand an infinite number of interaction terms and degrees of freedom. Thus, one needs an {\\it organizing principle} to make sensible calculations.  As a conventional way of classifying and organizing the interaction terms in our EFT approach as well as a guideline for the strengths of the various couplings the concept of naturalness has been adopted.    \\subsection{Naturalness}  We take an interaction Lagrangian, which involves the $\\sigma$ isoscalar-scalar meson and the $\\omega$ isoscalar-vector meson fields coupled to the nucleon field, defined  in Serot \\& Walecka (1997) as \\begin{eqnarray} \\!\\!\\!\\!\\! {\\cal L} =  \\sum_{i,k} \\frac{c_{i,k}}{i!k!} (\\frac{\\sigma}{f_{\\pi}})^{i} (\\frac{\\omega}{f_{\\pi}})^{k} \\left( \\frac{\\partial}{M}\\right) (\\frac{\\bar{\\psi}\\Gamma \\psi}{f^2_{\\pi}M})^{\\ell} f_{\\pi}^2 \\Lambda^{2}  \\, ,  \\label{1} \\end{eqnarray} with unknown expansion coefficients $c_{i,k}$ (overall coupling constants). In the expression above, $\\psi$ represents a baryon field, $\\Gamma$ is a Dirac matrix, and derivatives are denoted by $\\partial$. The resulting expressions of the Lagrangian density also contain factorial counting factors, needed to accomplish normalization. A direct generalization of this expression may involve additional fields such as  the $\\pi$, $\\varrho$ and $\\delta$ meson fields as well as the photon.  There is evidence from studies on properties of ordinary nuclear matter, that the expansion in the nonlinear meson coupling sector quickly converges. In fact, keeping only the cubic and quartic couplings of the $\\sigma$ meson one obtains, at saturation and intermediate densities, a semi-quantitative fit of nuclear matter properties (Boguta \\& Bodmer ~1997; Serot \\& Walecka ~1986). Nevertheless, a controlled expansion of the Lagrangian density to significantly higher densities ($\\rho \\geq 5 \\rho_0$), as those found in the interior of neutron stars, requires some assumptions on a natural ordering of the expansion coefficients. To accomplish that goal, we may express equation (\\ref{1}) in the form \\begin{eqnarray} \\!\\!\\!\\!\\!\\! \\! \\! {\\cal L} \\! = \\!  \\left( \\frac{\\partial \\, or \\, m_{\\pi}}{M}\\right)\\!\\! (\\frac{\\bar{\\psi}\\Gamma \\psi}{f^2_{\\pi}M})^{\\ell} f_{\\pi}^2 \\Lambda^{2} \\! \\sum_{i,k} \\frac{{\\tilde{c}}_{i,k}}{i!k!} (\\frac{g_{\\sigma} \\sigma}{M})^{i} (\\frac{g_{\\omega} \\omega}{M})^{k} \\, ,   \\label{2} \\end{eqnarray} where, in this expression, the Goldberger-Treiman relation ($g_s f_{\\pi} \\sim M$) was used to eliminate the $f_{\\pi}$ factor in the sum. The assumption of naturalness in the strong interaction physics means that, unless a more detailed explanation exists, all conceivable action terms that preserve the required fundamental symmetries should appear in the effective action of a theory with natural coupling coefficients  (Georgi ~1993; Manohar ~1996). Thus, the naturalness condition, when applied to an effective theory of the strong interaction, establishes that once the appropriate dimensional scales have been extracted using the {\\it naive dimensional analysis} proposed by Georgi (1993) and Manohar (1996),  the remaining dimensionless coefficients appearing in the effective action should all remain of order unity.  If the naturalness assumption is valid, then the effective strong interaction Lagrangian density can be truncated, with  an acceptable confidence, within the phenomenological physical domain of the theory. The  {\\it naive dimensional analysis} when applied in the formulation of an effective Lagrangian density involving nucleons and strongly interacting meson fields may be synthesized as follows: the amplitude of each strongly interacting field in the lagrangian, i.e. the meson fields, becomes dimensionless when divided by the pion decay weak constant; furthermore, to obtain the correct dimension ((energy)$^4$) for the Lagrangian density, an overall normalization scale $f_{\\pi}^2 \\Lambda^2 \\simeq f_{\\pi}^2 M^2$, with $M$ denoting the nucleon mass, has to be included; finally, for identical meson fields, self-interacting terms of power $n$ and a symmetrization factor $n!$ (for proper counting) should be included in the formalism. The overall dimensionless coefficients, after the dimensional factors and appropriate counting factors are extracted,  are of order ${\\cal O}(1)$ if naturalness holds. Of course, there is no general proof of the naturalness property, since no one knows how to derive the effective strong interaction Lagrangian density from Quantum Chromodynamics (QCD). Nevertheless,  the validity of naturalness and {\\it naive power counting rules} is supported by phenomenological studies.  At least two schemes allow a compact summation of the Lagrangian density (\\ref{2}). For $c_{i,\\kappa} = 1$, we have \\begin{eqnarray} \\!\\!\\!\\!\\!\\! {\\cal L} \\! = \\! \\left(\\frac{\\partial \\, or \\, m_{\\pi}}{M} \\right) \\!\\! \\left( \\frac{\\bar{\\psi} \\Gamma \\psi}{f_{\\pi}^2 M} \\right)^{\\ell} \\exp \\left( \\frac{\\sigma}{M} + \\frac{\\omega}{M} \\right)  f_{\\pi}^2 \\Lambda_{\\chi}^2 \\, ; \\end{eqnarray} and for $ c_{i,\\kappa} = i! \\kappa! $ we obtain \\begin{eqnarray} \\!\\!\\!\\!\\!\\! {\\cal L} = \\! \\left(\\frac{\\partial \\, or \\, m_{\\pi}}{M} \\right)  \\!\\! \\left( \\frac{\\bar{\\psi} \\Gamma \\psi}{f_{\\pi}^2 M} \\right)^{\\ell} \\!\\! \\left( \\frac{1}{\\!1 + \\frac{\\sigma}{M}} \\! \\right) \\! \\left( \\! \\frac{1}{1 + \\frac{\\omega}{M}} \\! \\right) f_{\\pi}^2 \\Lambda_{\\chi}^2 \\, . \\end{eqnarray}  \\subsection{Lagrangian Density}    Our strategy in this work is to consider a phenomenological and more flexible parametrization of the QHD Lagrangian density which combines the two previous limits and an extension of the interaction phase space of baryon and meson fields. The complete expression of our QHD interaction Lagrangian exhausts the whole fundamental baryon octet ($n$, $p$, $\\Sigma^-$, $\\Sigma^0$, $\\Sigma^+$, $\\Lambda$, $\\Xi^-$, $\\Xi^0$) and includes many-body forces simulated by nonlinear self-couplings and meson-meson interaction terms involving scalar-isoscalar ($\\sigma$, $\\sigma^*$), vector-isoscalar ($\\omega$, $\\phi$), vector-isovector ($\\mbox{\\boldmath$\\varrho$}$) and scalar-isovector ($\\mbox{\\boldmath$\\delta$}$).   The interaction term of our Lagrangian density, ${\\cal L}_{int}$, is defined as \\begin{eqnarray} & \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! \\sum_{B} \\!  \\bar{\\psi}_{B} \\! \\left[i \\gamma_{\\mu} \\partial^{\\mu} \\! - \\! \\gamma_{\\mu}\\Sigma^{\\mu}_{B \\, \\xi}\\!(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\omega) \\! - \\!  \\gamma_{\\mu} \\mbox{\\boldmath$\\tau$} \\! \\cdot \\! \\mbox{\\boldmath$\\Sigma$}_{B \\kappa}^{\\mu}\\!(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\mbox{\\boldmath$\\varrho$}) \\right. & \\nonumber \\\\ & \\left. \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! \\!\\!\\!\\!\\!- \\gamma_{\\mu}\\Sigma_{B \\eta }^{\\mu}(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\phi) \\! - \\! \\Sigma^{s}_{B \\varsigma}(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$})  \\right] \\!  \\psi_{B} & \\end{eqnarray} where one can identify the following self-energy insertions: \\begin{eqnarray} \\!\\!\\!\\!\\!\\!\\!\\!\\! \\Sigma^{\\mu}_{B \\, \\xi} (\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\omega) & \\! = \\! & g^*_{\\omega B \\xi}      \\omega^{\\mu} \\, \\, ; \\, \\mbox{\\boldmath$\\Sigma$}_{B \\kappa }^{\\mu}(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\mbox{\\boldmath$\\varrho$})  = \\frac{1}{2} g^*_{\\varrho B \\kappa } \\mbox{\\boldmath$\\varrho$}^\\mu \\, \\nonumber \\\\ \\!\\!\\!\\!\\!\\!\\!\\!\\! \\Sigma_{B \\eta }^{\\mu}(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}, \\phi) & \\! - \\! &  g^*_{\\phi B \\eta} \\phi^\\mu \\,  \\, ; \\, \\Sigma^{s}_{B \\varsigma}(\\sigma, \\sigma^*\\!\\!\\!, \\mbox{\\boldmath$\\delta$}) = M   \\Sigma^{s}_{B \\varsigma}  \\, . \\end{eqnarray} In these expressions, $g^*_{\\Phi}$, with $\\Phi =\\sigma, \\sigma^*, \\omega,  \\phi, \\varrho, \\delta$, re\\-pre\\-sents the effective parametrized baryon-meson coupling constants, defined as $g_{\\omega B \\xi}^* =  g_{\\omega B} m^*_{B \\xi}$, $g_{\\varrho B \\kappa}^* = g_{\\varrho B} m^*_{B \\kappa}$, $g_{\\phi B \\eta}^*  = g_{\\phi B} m^*_{B \\eta}$, and  $\\Sigma^{s}_{B \\varsigma} = m^*_{B \\varsigma}$, with \\begin{equation} m^*_{B \\alpha} \\equiv  \\left( 1 + \\frac{g_{\\sigma B} \\sigma + g_{\\sigma^* B} \\sigma^*+ \\frac{1}{2}g_{\\delta B} \\mbox{\\boldmath$\\tau$} \\cdot \\mbox{\\boldmath$\\delta$} }{\\alpha M_B} \\right)^{-\\alpha} \\, ; \\end{equation} ($\\alpha = \\xi$, $\\kappa$, $\\eta$, $\\varsigma$) (Taurines et al. 2001; Vasconcellos et al. 2012). Properties of the fields considered in our formulation are presented in table (\\ref{campos})). \\begin{center} \\begin{table}[t] \\vspace{-0.55cm} \\small \\caption{Properties of the fields considered in the formulation (\\ref{NLD}). In what follows, we use the abbreviations: ISS: isoscalar-scalar; IVS: isovector-scalar; ISV: isoscalar-vector; IVV: isovector-vector.} \\begin{center} \\begin{tabular}{llclc} \\hline Fields & Classification & Particles & Coupling   & Mass \\\\ & & & Constants & (MeV) \\\\ \\hline $\\psi_B$ & Baryons & N, \\,$\\Lambda$,& N/A & 939, 1116, \\\\ &           &              $\\Sigma$, \\,$\\Xi$   &     &       1193, 1318\\\\ $\\psi_l$ & Leptons & $e^-$, $\\mu^-$& N/A &0,5,\\,106 \\\\ \\hline $\\sigma $ & ISS-meson  & $\\sigma $ & $g^*_{\\sigma_B}$ & 550 \\\\ $\\mbox{\\boldmath$\\delta$}$ & IVS-meson  & $a_0$ &$g^*_{\\delta_B}$ &980\\\\ $\\omega_\\mu $ & ISV-meson  & $\\omega $ &$g^*_{\\omega_B}$ &782 \\\\ $\\mbox{\\boldmath$\\varrho$}_{\\mu}$ & IVV-meson  & $\\rho$ & $g^*_{\\varrho_B}$&770 \\\\ $\\sigma^\\ast $ & ISS-meson  & $f_0$ &$g^*_{\\sigma\\ast_B}$ & 975 \\\\ $\\phi_\\mu $ & ISV-meson & $\\phi $ &$g^*_{\\phi_B}$ & 1020 \\\\ \\hline \\end{tabular} \\label{campos} \\end{center} \\end{table} \\end{center}      The complete expression for the Lagrangian density in the mean field approximation, $ {\\cal L}_{\\xi \\varsigma  \\kappa \\eta }$, becomes \\begin{eqnarray} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!  & & \\frac{1}{2}m_\\sigma^2\\sigma_0^2 + \\frac{1}{2} m_{\\sigma\\ast}^2 \\sigma_0^{\\ast 2} +\\frac{1}{2}m_\\omega^2\\omega_0^2 +\\frac{1}{2} m_\\phi^2 \\phi_0^2 \\\\ & \\! + \\! & \\frac{1}{2} m_\\rho^2 \\varrho_{03}^2 + \\frac{1}{2} m_\\delta^2 \\delta_{3}^2   + \\sum_l \\bar\\psi_l \\,(i \\gamma_\\mu \\partial^\\mu - m_l)\\,\\psi_l \\nonumber \\\\  & \\! + \\! & \\sum_B \\bar{\\psi}_B (i \\gamma_\\mu \\partial^\\mu -  g_{\\omega B} m^*_{B \\xi}\\gamma_0 \\omega_0   -   M_{B \\varsigma}^{\\ast}  ) \\psi_B \\nonumber \\\\  &  \\! - \\! & \\sum_B  \\! \\bar{\\psi}_B  ( \\frac{1}{2} g_{\\varrho B} m^*_{B  \\kappa} \\gamma_0 \\tau^{(3)} \\varrho_{03}   \\! + \\! g_{\\phi B} m^*_{B  \\eta} \\! \\gamma_0  \\phi_0 ) \\psi_B \\, , \\nonumber\\label{NLD} \\end{eqnarray} with the effective baryon mass defined as $ M_{B \\varsigma}^{\\ast} =  M_B m^*_{B \\varsigma}$.  The resulting expression for the Lagrangian density is known as the parameterized coupling model. Among the numerous possibilities of choice of parameters for the model we may focus our attention in variations of the $\\varsigma$ pa\\-ra\\-me\\-ter keeping $\\xi = \\kappa = \\eta = 0$ (scalar or S-model).  This parametrization reduces to the Walecka's QHD-I model (Serot \\& Walecka 1997) if $\\varsigma = 0$ and $g_{\\varrho B} = g_{\\phi B} = 0$. In case only $g_{\\varrho B} \\neq 0$, this parametrization reduces to the Walecka's QHD-II model (Serot \\& Walecka 1997). Assuming $ \\varsigma = 1$, performing a binomial expansion of $m_{B \\varsigma}$, truncating the perturbative series to cubic and quartic self-interactions terms involving the scalar-isoscalar $\\sigma$ meson (making $g_{\\sigma^* B} = g_{\\delta B} = 0$), this parametrization reduces to the model of Boguta \\& Bodmer (1997). The second choice contemplates variations of the parameters $\\varsigma$ and $\\xi$ while keeping $\\kappa = \\eta = 0$ (scalar-isoscalar-vector or SISV-model). In the third choice we may consider variations of $\\varsigma$, $\\xi$, and $\\kappa$, and fixing $\\eta = 0$ (scalar-isoscalar-vector-isovector-vector or SISVIVV-I-model), as already discussed in (Dexheimer, Vasconcellos \\& Bodmann ~2008). Finally, in the fourth choice we may consider variations of the four parameters of the theory, $\\varsigma$, $\\xi$, $\\kappa$, and $\\varsigma$ (scalar-isoscalar-vector-isovector-vector  or SISVIVV-II-model). These examples of parameterizations of our approach are shown in Table (\\ref{parameterizations}). \\begin{center} \\begin{table}[htdp] \\caption{Examples of parameterizations of our model. S: scalar model; SISV: scalar-isoscalar-vector model; SISVIVV: scalar-isoscalar-vector-isovector-vector  model. Model II differs from I, due to the presence of the $\\phi$ meson.} \\centerline{\\begin{tabular}{ccccc} \\hline Model & $\\varsigma$ & $\\xi$ & $\\kappa$ & $\\eta$ \\\\ \\hline S & $ \\neq 0$ & $0$ & $0$ & $0$ \\\\ \\hline SISV & $\\neq 0$ & $\\neq 0$ & $0$ & $0$ \\\\ \\hline SISVIVV-I & $\\neq 0$ & $\\neq 0$ & $\\neq 0$  & $0 $ \\\\ \\hline SISVIVV-II & $\\neq 0$ & $\\neq 0$ &$\\neq 0$ &$\\neq 0$ \\\\ \\hline \\end{tabular} } \\label{parameterizations} \\end{table} \\end{center}  We consider, as an example, the expansion of the term $g_{\\omega B} m^*_{B \\xi}\\gamma_0 \\omega_0$ in expression (\\ref{NLD}), in the particular case $\\xi=1$. In the framework of a local mean field approximation, expectation values of meson fields correspond to classical numbers\\footnote{It is well known that in the realm of the mean field approximation, expectation values of meson fields may be treated as classical numbers in space-time despite the density dependence of these fields in Fermi space.}. In this context, taking into account that $\\frac{g_{\\sigma B} \\sigma_0 + g_{\\sigma^* B} \\sigma^*_0+ \\frac{1}{2}g_{\\delta B} \\tau_3 \\delta_{30} }{M_B} << 1$, we may use the binomial theorem to expand the contribution $g_{\\omega B} m^*_{B \\xi}\\gamma_0 \\omega_0$ for the particular choice $\\xi = 1$: \\begin{eqnarray} & g_{\\omega B} m^*_{B 1}\\gamma^0 \\omega_0   \\sim  g_{\\omega B}\\gamma^0 \\omega_0 & \\nonumber \\\\ & -  g_{\\omega B} \\left( \\frac{g_{\\sigma B} \\sigma_0 + g_{\\sigma^* B} \\sigma^*_0+ \\frac{1}{2}g_{\\delta B} \\tau_3 \\delta_{30} }{M_B}\\right)\\gamma^0 \\omega_0 & \\nonumber \\\\ & +  g_{\\omega B} \\left( \\frac{g_{\\sigma B} \\sigma_0 + g_{\\sigma^* B} \\sigma^*_0+ \\frac{1}{2}g_{\\delta B} \\tau_3 \\delta_{30} }{M_B}\\right)^2 \\gamma^0 \\omega_0 & \\nonumber \\\\ & -   g_{\\omega B} \\left( \\frac{g_{\\sigma B} \\sigma_0 + g_{\\sigma^* B} \\sigma^*_0+ \\frac{1}{2}g_{\\delta B} \\tau_3 \\delta_{30} }{M_B}\\right)^3 \\gamma^0 \\omega_0  ... & \\end{eqnarray} From this expression we can identify many-body attractive and/or repulsive coupling terms like for instance $\\sigma_0 \\omega_0$, $\\sigma_0^2 \\omega_0$, $\\sigma^{*2}_0 \\omega_0$, $\\delta_{30} \\omega_0$, $\\delta_{30}^2 \\omega_0$, $\\sigma_0  \\sigma^{*} \\delta_{30} \\omega_0$ and many others. Similarly, the remaining terms of the Lagrangian density also exhibit additional many-body coupling terms. Again, the final selection of the contributions to be considered requires an analysis of the formal coherence of the theory. In this sense, it is important to remember that the theory must embody fundamental symmetries, conservation laws as well as physical properties which are relevant for strong in\\-te\\-racting relativistic nuclear many-body systems, such as Lorentz covariance, microscopic causality, naturalness, analyticity, uniqueness, among others.   From expression (\\ref{1}), an extension of the model of Serot \\& Walecka (1997), which exhausts the whole fundamental baryon octet ($n$, $p$, $\\Sigma^-$, $\\Sigma^0$, $\\Sigma^+$, $\\Lambda$, $\\Xi^-$, $\\Xi^0$) and includes baryon-meson interaction terms involving scalar-isoscalar ($\\sigma$, $\\sigma^*$), vector-isoscalar ($\\omega$, $\\phi$), vector-isovector ($\\varrho$) and scalar-isovector ($\\delta$) may be synthesized as \\begin{eqnarray} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! {\\cal L} & \\!\\! = \\!\\! & \\sum_B \\sum_{\\wp} \\! \\frac{c_{\\wp}}{\\Pi_{\\wp} \\wp! } \\left( \\frac{\\partial \\, or \\, m_{\\pi}}{M_B}\\right) (\\frac{\\bar{\\psi}_B \\Gamma \\psi_B}{f^2_{\\pi}M_B})^{\\ell} f_{\\pi}^2 \\Lambda^{2} \\nonumber \\\\ & \\times & (\\frac{\\sigma}{f_{\\pi}})^{i} (\\frac{\\sigma^*}{f_{\\pi}})^{j} (\\frac{\\omega}{f_{\\pi}})^{k} (\\frac{\\varrho}{f_{\\pi}})^{m} (\\frac{\\delta}{f_{\\pi}})^{n} (\\frac{\\phi}{f_{\\pi}})^{q} \\, , \\label{extension} \\end{eqnarray} with $\\wp = i, j, k, m, n, q$. The {\\it natural} limit of this expression, namely, by assuming $c(\\wp) \\to 1$, is \\begin{eqnarray} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!  {\\cal L}   =  \\sum_B \\left(\\frac{\\partial \\, or \\, m_{\\pi}}{M_B} \\right) \\!\\! (\\frac{\\bar{\\psi}_B \\Gamma \\psi_B}{f^2_{\\pi}M_B})^{\\ell}  f_{\\pi}^2 \\Lambda^{2} \\exp \\sum_{\\iota}\\frac{g_{\\iota} \\Phi_{\\iota}}{M_B} \\, , \\end{eqnarray} with $\\Phi_{\\iota}=\\sigma, \\sigma^*, \\omega,  \\phi, \\varrho, \\delta$. Through systematic choices for the parameters of the Lagrangian density (\\ref{NLD}) we also obtain, in the limit of naturalness, exponential couplings. ",
        "conclusions": "In the following, we determine the EoS and population profiles for neutron stars assuming local charge neutrality and beta equilibrium for the scalar version of our approach. By solving the Tolman-Oppenheimer-Volkoff (TOV) equations (Tolman 1939; Oppenheimer \\& Volkoff ~1939) we obtain the mass-radius relation for families of neutron stars with hyperon content. Our results for different values of the $\\lambda$ parameter are illustrated in Figs. (\\ref{eos}), (\\ref{popl006}), (\\ref{popl010}), (\\ref{popl014}), and (\\ref{tov}). \\begin{table}[h,t,b] \\vspace{-0.4cm} \\caption{\\label{cc} Coupling parameters of our approach.} \\begin{tabular}{cccccc} \\small $\\lambda$ & $(\\frac{M^*}{M})_{_0}$ &  $K_0$~(MeV) & $(\\frac{g_{\\sigma B}}{m_{\\sigma}})^2$ & $(\\frac{g_{\\omega B}}{m_{\\omega}})^2$ & $(\\frac{g_{\\varrho B}}{m_{\\varrho}})^2$ \\\\ \\hline & &  & & & \\\\ 0.06 & 0.70  & 262   &  11.87  & 6.49  & 3.69  \\\\ 0.10 & 0.75  & 226   &  10.42  & 5.10  & 3.94  \\\\ 0.14 & 0.78  & 216   &  9.51   & 4.32   & 4.07 \\\\ \\end{tabular} \\end{table} In the figures, each value of $\\lambda$  generates a sequence of neutrons stars, where each star will have a different EoS, particle population, central density, and maximum mass. The parametrization that generates the higher pressure and maximum mass of a neutron star is the one with the lower possible value of the parameter $\\lambda$, ie $\\lambda = 0.06$, which provides a mass of $1.97\\,M_{\\odot}$, in very good agreement with recent observations (Demorest et al. ~2010). This result corresponds to the parametrization which provides a smaller effective baryon mass and, hence, a more bound matter. Although smaller lambda values originate higher values for the maximum mass of neutron stars, this parameter should attain a minimum value that allows the model to reproduce nuclear saturation properties, like for example a compressibility modulus smaller than 300 MeV.      The analysis of these results demands first to remember that a stiffer, or equivalently, more rigid equation of state of nuclear matter is related to higher values of the internal pressure of the system and, accordingly, to higher values of the compressibility modulus $|K_{sym}|$ of nuclear matter. This in turn requires stronger contributions from repulsive components of the nuclear force when compared to the attractive ones. In our general approach, however, many body forces lowers the intensities both of attractive and repulsive interaction terms due to {\\it shielding effects}, which result in higher (lower) values of the compressibility modulus $|K_{sym}|$ of nuclear matter in the case of higher (lower) relative reduction of the attractive (repulsive) contributions.   Our present results (scalar model) are consistent with the analysis above, since we have obtained a stiffer equation of state for lower values of the $\\lambda$ parameter ($\\lambda = 0.06$), which corresponds to the stronger relative reduction of the attractive contribution (see Eq. (\\ref{effectivemass})). Si\\-mi\\-lar\\-ly, the results for the particle po\\-pu\\-la\\-tion show that the threshold equation for a given species (Glendenning ~1996), $$\\mu_n \\! - \\! q_B \\mu_e \\! - \\! g_{\\omega B}\\omega_0 \\! - \\! g_{\\varrho B} \\varrho_{0 3} I_{3 B} \\geq M_B ( 1 \\! - \\! g_{\\sigma B} \\sigma_0/M_B),$$ is also affected by the presence of many-body forces components in the strong interaction sector (sse Eq. (\\ref{effectivemass}): \\begin{eqnarray} & \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! \\mu_n  \\! - \\! q_B \\mu_e  \\! - \\! g_{\\omega B}\\omega_0 \\! - \\! g_{\\varrho B} \\varrho_{0 3} \\! I_{3 B} \\! \\geq \\! M_B ( 1 \\! - \\! g^*_{\\sigma B} \\sigma_0/M_B) & \\nonumber \\\\ & \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!  \\simeq  \\! M_B \\! \\left( \\! 1 - \\frac{g_{\\sigma B}\\sigma_0}{M_B} + \\left( \\! \\begin{array}{c} \\lambda \\\\ 2  \\end{array} \\right) \\! \\left( \\! \\frac{g_{\\sigma B} \\sigma_0}{\\lambda M_B} \\! \\right)^2   \\right) + {\\cal O}(3) \\, , & \\label{TE} \\end{eqnarray} where $\\mu_n$ and $\\mu_e$ represent respectively the neutron and electron chemical potential, and $q_B$ is the baryon charge. According to Eq. (16) many body forces shift the critical density for hyperon sa\\-tu\\-ra\\-tion to higher densities and originate an anti-correlation between the amount of hyperons and the values of $\\lambda$: to the lowest va\\-lues of $\\lambda$ cor\\-res\\-pond larger amounts of hyperons in the system, which may cause the softening of the EoS. However, our results show that the shielding effect of the long-range nuclear components of the nuclear force is predominant. The first hyperon species that appears is the $\\Sigma^{-}$, closely followed by the $\\Lambda$, since the negative charge of the $\\Sigma^{-}$ outweighs the 80 MeV mass difference, as a result of the more lenient conditions $\\mu_{\\Lambda} = \\mu_n$ and $\\mu_{\\Sigma^-} = \\mu_n + \\mu_e$. The $\\Sigma^{-}$ fraction saturates at about $0.3$, while $\\Lambda$, free of isospin-dependent forces, continue to accumulate until short-range repulsion forces cause them to saturate. Other hyperon species follow at higher densities.  \\begin{figure}[h,t,b] \\vspace{-0.4cm} \\centering \\includegraphics[width=8.0cm]{eos.eps} \\caption{\\label{eos} Dependence of the equation of state (EoS) on $\\lambda$.} \\end{figure} \\begin{figure}[h,t,b] \\centering \\includegraphics[width=8.0cm]{pop_l006.eps} \\caption{Particle population for $\\lambda=0.06$. In the figure, the labels $\\rho_B$ and $\\rho_0$ indicate respectively baryon density and nuclear saturation density.  \\label{popl006}} \\end{figure} \\begin{figure}[h,t,b] \\vspace{-0.4cm} \\centering \\includegraphics[width=8.0cm]{pop_l010.eps} \\caption{Particle population for $\\lambda=0.10$. The symbols have the same meaning as in the previous figure.\\label{popl010}} \\end{figure} \\begin{figure}[h,t,] \\centering \\includegraphics[width=8.0cm]{pop_l014.eps} \\caption{Particle population for $\\lambda=0.14$. Symbols have the same meaning as in previous figures. \\label{popl014}} \\end{figure} \\begin{figure}[h,t] \\vspace{-0.4cm} \\centering \\includegraphics[width=8.0cm]{tov.eps} \\caption{Mass-radius diagram for three different parameterizations of the scalar version of the model: $\\lambda = 0.06,\\,0.10,\\,0.14$. \\label{tov}} \\end{figure}   It is our understanding that the naturalness condition is equivalent to exhaust the phase space of the fundamental interactions. This depletion of the phase space can be accomplished by including in nuclear matter the largest possible number of {\\it information}, ie. baryon and meson degrees of freedom and additionally considering many body forces between baryon and meson fields, as well as self-coupling terms involving meson fields. When the condition of na\\-tu\\-ral\\-ness is achieved, it is unnecessary to get rid of hea\\-vier degrees of freedom by integrating then out. In this case, the effects of heavier degrees of freedom are not anymore implicitly contained in coupling parameters of the effective field theory. An interesting aspect of our study is that appa\\-ren\\-tly, many body forces occupy a larger role than originally thought in the description of properties of nuclear system at high density. Moreover, the condition of naturalness is achieved with only a small part of the parameter set. This result was expected and can be explained by the saturation property of nuclear matter: for parameter values less than $3$, the model with parameterized couplings completely exhausts the phase space of many-body interactions.    The model with parameterized couplings shows promising results. However, the model needs to broaden its scope by the adoption of new parameterizations, expanding the set of parameters of the theory. Interesting issues for future studies will also be the role of finite temperature, neutrino trapping and strong magnetic effects in neutron stars. Work along these lines is in progress.   \\vspace{-0.35cm}"
    },
    "1402/1402.5412_arXiv.txt": {
        "abstract": "Adopting our new method for matching general relativistic, ideal magnetohydrodynamics to its force-free limit, we perform the first systematic simulations of force-free pulsar magnetospheres in general relativity. We endow the neutron star with a general relativistic dipole magnetic field, model the interior with ideal magnetohydrodynamics, and adopt force-free electrodynamics in the exterior. Comparing  the spin-down luminosity to its corresponding Minkowski value, we find that general relativistic effects give rise to a modest enhancement: the maximum enhancement for $n=1$ polytropes is $\\sim 23\\%$. Evolving a rapidly rotating $n=0.5$ polytrope we find an even greater enhancement of $\\sim 35\\%$. Using our simulation data, we derive fitting formulas for the pulsar spin-down luminosity as a function of the neutron star compaction, angular speed, and dipole magnetic moment. We expect stiffer equations of state and more rapidly spinning neutron stars to lead to even larger enhancements in the spin-down luminosity. ",
        "introduction": "\\label{sec:introduction}  Pulsars are believed to be rapidly rotating neutron stars (NSs) that lose their rotational kinetic energy primarily due to emission of electromagnetic radiation (see e.g. ~\\cite{ShaTeu83a}). Pulsars are extremely accurate clocks that can be used to probe fundamental physics, such as the nuclear equation of state (EOS), and theories of gravity (see e.g. \\cite{lrr-2008-9}). They can even function as detectors of gravitational waves~\\cite{0264-9381-27-8-084013}.  Pulsars are detectable from radio to gamma-ray frequencies. However, only a small number of pulsars is observable in the visible bands \\cite{2002nsps.conf...44S}. Fermi-LAT technology \\cite{Atwood:2009ez} has allowed the recent discovery of numerous gamma-ray pulsars. To date there are over 2300 known pulsars \\cite{2005AJ....129.1993M} of which over $130$ are gamma-ray pulsars \\cite{FERMILATGPULSARS,TheFermi-LAT:2013ssa}: the vast majority of pulsars are radio pulsars. The most slowly rotating  pulsar is PSR-J2144-3933~\\cite{2005ApJ...627..397Z}, with a period of 9.43 s. The most rapidly rotating pulsar is PSR-J1748-2446ad \\cite{2006Sci...311.1901H}, with a period of 1.395 ms. Hence pulsar periods cover almost 4 orders of magnitude.  A complete theory that can explain both the radio and the gamma-ray emission remains elusive.  Nevertheless, our understanding of pulsar physics has been drastically improved in recent years as a result of computational simulations that model the global pulsar magnetosphere.  Rapid rotation in pulsars induces strong electric fields capable of stripping matter off the NS surface and eventually populating the exterior with tenuous plasma \\cite{Goldreich:1969sb}. Motivated by this result, most global studies of pulsar magnetospheres adopt the force-free (FFE) limit of magnetohydrodynamics (MHD), which is valid in such environments. Moreover, to simplify the analysis further, most studies consider the case where the NS dipole magnetic moment is aligned or antialigned with its angular momentum. This simple model has been successful in producing the main features of the pulsar magnetosphere of an aligned rotator: (i) a near or closed zone, which corotates with the star, and in which magnetic field lines return to the stellar surface; (ii) a far or open zone, which extends beyond the light cylinder radius $R_{\\rm LC}$, and in which the magnetic lines are open and extend to infinity; (iii) a Y point at the location of the light cylinder, at which the magnetic field lines first open; and (iv) an equatorial current sheet. All these features are predictions of the so-called pulsar equation \\cite{Michel73,Scharleman73}.  The first successful numerical solution of the pulsar equation was presented in~\\cite{Contopoulos:99}, which was later followed by numerous studies ~\\cite{2002astro.ph.11141K,McKinney:2006sd, 2006ApJ...648L..51S,Gruzinov:2008gb,Gruzinov:2013pva, Tchekhovskoy:2012hm,Uzdensky:2012tf} that probed the global features of an aligned rotator in flat spacetime and reinforced the global picture outlined above. These studies did not include the magnetized NS interior and modeled the effects of rotation through a boundary condition on the spherical stellar surface. Simulations of force-free magnetospheres have produced important results, such as a proof of existence of a stationary force-free magnetospheric configuration, the calculation of the spin-down luminosity of force-free oblique rotators \\cite{2006ApJ...648L..51S}, and the evolution of the obliquity angle \\cite{Philippov:2013aha}, all in flat spacetime. Recently, there have also been some analytic efforts to understand the emission from an accelerated isolated pulsar in flat spacetime (see e.g. \\cite{Brennan:2013ppa,Gralla:2014yja}).  In addition to assuming that pulsars possess dipole magnetic fields, the common assumptions and simplifications in these  earlier studies have been the following: \\begin{enumerate} \\item \\label{ass1} The pulsar magnetosphere is well described by force-free electrodynamics.  \\item \\label{ass2} The backreaction of the magnetic field onto the interior matter is ignored.  \\item \\label{ass3} Deviations from sphericity of the stellar surface (e.g. due to rapid rotation) are ignored.  \\item \\label{ass4} Curved spacetime effects are ignored. \\end{enumerate}  For determining the pulsar spin-down luminosity, assumption \\ref{ass1} is further justified by new results reported in \\cite{Philippov:2013tpa}, where a particle-in-cell simulation of the global pulsar magnetosphere yields a spin-down luminosity consistent with earlier force-free studies~\\cite{2006ApJ...648L..51S}.  While assumption \\ref{ass2} is likely to break down in a thin layer near the stellar surface~\\cite{Mestel94,Contopoulos:1999ga,Morrison:2004fp}, the calculations reported in ~\\cite{Palenzuela:2012my} indicate that the backreaction of the magnetic field can lead only to small corrections to the spin-down luminosity. To our knowledge the extent to which assumption \\ref{ass3} is correct has never been studied before (likely due to the complexity of imposing boundary conditions on a nonspherical stellar surface). Relaxing assumption \\ref{ass4} has been argued to lead to important effects.  For example, \\cite{Deutsch55} pointed out that general relativity (GR) can amplify the induced electric field due to the rotation of the star.  It has also been argued \\cite{Beskin90,Musmilov92} that frame dragging induces an enhanced electric field which contributes to particle acceleration in the polar cap region.  More recently, the {\\it vacuum} Maxwell equations were solved in the curved spacetime of a {\\it slowly} rotating NS endowed with a dipole magnetic field~\\cite{Petri:2014oza}. It was shown that GR results in an enhancement in the spin-down luminosity of a pulsar in vacuum up to $\\sim 34\\%$ (see also \\cite{Rezzolla:2004hy} for an earlier analytic study). Furthermore, adopting a GR resistive MHD scheme, it has been reported that the spin-down luminosity from a GR aligned rotator deviates from the corresponding flat spacetime value by $\\sim 20\\%$~\\cite{Palenzuela:2012my}.  As noted in ~\\cite{Palenzuela:2012my} this deviation could be due to the adopted resistive MHD scheme, to the way the flat spacetime spin-down luminosity formula is applied when dealing with oblate stars or to GR effects. The origin of the difference could also be due to the inclusion of backreaction of the magnetic fields onto the thin layer near the NS surface or some combination of all these factors. Finally, force-free simulations of rotating and nonrotating, collapsing neutron stars have recently been performed in \\cite{Lehner:2011aa}. All these earlier studies indicate that GR plays an important role in pulsar magnetospheres.  In this paper we perform the first systematic study of general relativistic effects in the force-free pulsar magnetospheres of aligned rotators. We adopt ideal GRMHD for the NS interior, and model its magnetosphere by adopting general relativistic force-free electrodynamics. Using the numerical method we presented and tested in~\\cite{Paschalidis:2013gma} we match the ideal MHD dense interior to the force-free exterior. We evolve the stars until the systems relax and compute the spin-down luminosity after steady state has been achieved.  As in most previous studies we neglect the back-reaction of the magnetic field onto the matter interior. Consequently, all our results scale with the magnetic field strength (or the dipole magnetic moment).  We split our study into two stages: in the first stage we treat the ``slow'' rotation limit. In particular we generalize the flat spacetime results for the spin-down luminosity by considering sequences of Tolman-Oppenheimer-Volkoff (TOV) spherical stars, endowing these stars with a slow, uniform, rotational velocity, and a general relativistic dipole magnetic field~\\cite{Wasserman83}. When both $T/|W| \\ll 1$ and ${\\cal M}/|W| \\ll 1$, where $T$, ${\\cal M}$ and $|W|$ are the kinetic, magnetic and gravitational binding energies, respectively, the stellar structure remains unaffected and nearly spherical. These simulations allow us to isolate and quantify the effects of the {\\it compaction} on the pulsar spin-down luminosity, as frame dragging is not present. We find that the higher the compaction, the larger the enhancement in the spin-down luminosity when compared to its flat spacetime value. The maximum enhancement reaches $\\sim 9.3\\%$ near the maximum compaction limit.  In the second stage, we treat the ``rapid'' rotation limit by constructing self-consistent equilibrium sequences of uniformly rotating, polytropic NSs using the Cook-Shapiro-Teukolsky (CST) code~\\cite{Cook:1993qj,Cook:1993qr}.  Endowing these stars with a general relativistic dipole magnetic field (assuming ${\\cal M}/|W| \\ll 1$, hence leaving the nonspherical structure of the rotating star unaffected by the magnetic field) and evolving until relaxation, we find that rotation (frame dragging) increases the enhancement of the spin-down luminosity over the corresponding Minkowski value even further. The maximum enhancement for $n=1$ polytropes is $\\sim 23\\%$. Evolving a rapidly rotating $n=0.5$ polytrope we find an even greater enhancement of $\\sim 35\\%$. Using our simulation data, we derive fitting formulas for the pulsar spin-down luminosity as a function of the NS compaction, angular speed, and dipole magnetic moment. We argue that general relativistic effects are responsible for the observed enhancement and not the distorted surface of the star by considering the flat spacetime evolution of a highly oblate, magnetized star. We find that the spin-down luminosity in this case is practically the same as for a spherical star. We expect that stiffer equations of state and more rapidly rotating stars should lead to even larger enhancements in the pulsar spin-down luminosity.  This paper is organized as follows. In Sec.~\\ref{sec:numerical_setup} we briefly summarize our numerical method for evolving and matching ideal GRMHD to its force-free electrodynamics (GRFFE) limit. We also present the diagnostics we adopt to monitor these systems. A detailed description of our initial data is presented in Sec.~\\ref{sec:ID}. A set of tests and results are summarized in Sec.~\\ref{sec:results}, where we also include a resolution study. We conclude in Section~\\ref{sec:conclusions} with a summary  and a discussion of  astrophysical implications. Throughout we adopt geometrized units where $G=c=1$. ",
        "conclusions": "\\label{sec:conclusions}  Magnetized neutron stars possess a force-free magnetosphere.  Pulsar magnetospheres have been studied numerically over the last two decades. However, all of these early force-free studies were carried out in flat spacetime. It has been suggested that general relativistic effects may become important but due mainly to frame dragging~\\cite{Beskin90,Musmilov92}.  Using the new method we recently developed for matching general relativistic ideal MHD to its corresponding force-free limit, we have performed the first systematic study of aligned rotator force-free magnetospheres in GR. We constructed equilibrium sequences of ``slowly'' and rapidly rotating NSs. The former are modeled as incompressible TOV stars covering almost the entire allowed range of compaction. The latter are constant $\\Omega$ sequences of uniformly, rapidly rotating, GR, polytropic stars generated by the CST code, and ranging from the mass-shedding limit to the maximum compaction configuration for a given stellar angular frequency $\\Omega$. We endowed these stars with a weak GR dipole magnetic field and evolved the fields to relaxation.  Some of our rapidly rotating NS models (those close to the mass-shedding limit) are highly distorted (see Table \\ref{tab:self_rot_ID}). To ensure that any deviations from the flat spacetime spin-down luminosity are due to strong field effects and not to the distorted surface of the stars, we evolved a highly oblate star (with eccentricity $e=0.799$) in flat spacetime, and found that the outgoing Poynting luminosity is independent of the shape of the stellar surface, i.e., the resulting spin-down luminosity is approximately the same as if the star were a sphere spinning with the same angular frequency and endowed with the same magnetic field.  Normalizing the spin-down luminosity $L$ to its corresponding Minkowski value $L_0$, we find that GR effects give rise to a modest enhancement. As both the compaction and the stellar angular frequency increase, the enhancement becomes more pronounced. In the ``slow'' rotation limit, where we isolate the effects of compaction, the maximum enhancement is $\\sim 9.3\\%$, independent of the equation of state. However, for rapidly rotating stars, where frame dragging is also important, the maximum enhancement for $n=1$ polytropes is $\\sim 23\\%$, and for a rapidly rotating $n=0.5$ polytrope we find an enhancement of $\\sim 35\\%$.  We expect stiffer equations of state and more rapidly neutron rotating stars to lead to even larger enhancements in the spin-down luminosity. For the cases we studied here we provided fitting functions for the general relativistic spin-down luminosity of the form $L=G(Z_p,\\Omega)L_0$, where $Z_p$ is the gravitational redshift of light emitted from the NS pole as measured by a static observer at infinity.  Our results show that even moderate compaction stars (for which frame dragging is accounted for) have larger $L/L_0$ than the highest compaction TOV model (for which frame dragging is not accounted for). Hence, it is natural to conclude that there is a strong correlation between frame dragging and the enhancement of the spin-down luminosity over its value in flat spacetime. This suggestion is consistent with the findings of earlier theoretical arguments~\\cite{Beskin90,Musmilov92} of why GR pulsar magnetospheres are different than those in flat spacetime. Further studying the source of the  differences between GR and flat spacetime studies of  pulsars, is beyond the scope of this paper, and will be the subject of a future work.  Should future gravitational wave observations prove able to constrain the nuclear equation of state (see e.g. \\cite{Lackey:2011vz,Bauswein:2011tp,Paschalidis:2012ff} and references therein), our study of pulsar magnetospheres and similar future studies can constrain the NS magnetic field strength and geometry as follows: pulsars in detached, mildly relativistic binaries allow for a determination of the pulsar mass, the angular frequency of rotation, and the spin-down rate. If the nuclear equation of state is known then the compaction of the NS follows from the mass-radius relationship. As the spin-down luminosity depends on the compaction, the angular speed and the magnetic field, then from observations one can constrain the magnetic field, given that all the other parameters are in principle measurable or can be inferred. Even the obliquity angle between the magnetic dipole moment and the rotation axis could possibly be constrained, once oblique rotators are studied in GR. This motivates an investigation of such rotators in GR, which we intend to perform in a future study."
    },
    "1402/1402.5248_arXiv.txt": {
        "abstract": "In this paper, we show how the two-fluid equations describing the evolution of a dust and gas mixture can be re-formulated to describe a single fluid moving with the barycentric velocity of the mixture. This leads to evolution equations for the total density, momentum, the differential velocity between the dust and the gas phases and either the dust-to-gas ratio or the dust fraction. The equations are similar to the usual equations of gas dynamics, providing a convenient way to extend existing codes to simulate two-fluid mixtures without modifying the code architecture.  Our approach avoids the inherent difficulties related to the standard approach where the two phases are separate and coupled via a drag term. In particular, the requirements of infinite spatial and temporal resolution as the stopping time tends to zero are no longer necessary. This means that both small and large grains can be straightforwardly treated with the same method, with no need for complicated implicit schemes. Since there is only one resolution scale the method also avoids the problem of unphysical trapping of one fluid (e.g. dust) below the resolution of the other.  We also derive a simplified set of equations applicable to the case of strong drag/small grains, consisting of the standard fluid equations with a modified sound speed, plus an advection-diffusion equation for the dust-to-gas ratio. This provides a simple and fast way to evolve the mixture when the stopping time is smaller than the Courant timestep. We present a Smoothed Particle Hydrodynamics implementation in a companion paper. ",
        "introduction": "\\label{sec:intro} Dust. It pervades the interstellar medium, molecular clouds, young stellar systems and protoplanetary discs, and is the material from which planets are formed. Grain properties are inferred via a range observational techniques at multiple wavelengths, including spectral emission distributions, scattered light images, thermal emission maps, mid-infrared spectroscopy, polarimetry and molecular emissions (see \\citealt{Pinte2008} for a detailed analysis). We can thus determine the spatial distribution, composition, crystallinity and even more importantly, grain sizes $s$, The largest fraction of the dust mass lies in the small grains since the dust size distribution is found to vary as $n\\left( s \\right) \\simeq s^{-3.5}$ \\citep{Mathis1977}.  Simulating grain evolution in astrophysical objects is of primary importance since they play a crucial role in transferring energy, being the main source of thermal infrared emission. They thus affect the thermodynamics of the gas and --- in turn --- its dynamics. This is key to astrophysical processes that are controlled by the ability of the system to dissipate its gravitational energy, such as star and planet formation. However, computing the local opacities self-consistently taking into account the local evolution of the dust population and in general, simulating dust and gas mixtures evolution is quite challenging. Dust grains mainly interact with the surrounding gas by microscopic collisions between gas molecules and the dust grains. This results in a macroscopic drag force which damps the differential velocity between the two phases. The typical timescale for the damping is called the stopping time $t_{\\rm{s}}$, which usually increases when the grain size decreases. Numerical methods thus need to be efficient over a wide range of dust sizes encountered in astrophysical systems. This is particularly the case in protoplanetary discs, where grains with different sizes become spatially differentiated by dynamical processes \\citep{Weidendust1977,NSH86,GaraudLin2004,Youdin2005,Carballido2006}.  In two previous papers (\\citealt{LP12a,LP12b}, hereafter \\citetalias{LP12a,LP12b}), we have detailed and benchmarked a Lagrangian Smoothed Particle Hydrodynamics (SPH) formalism to treat astrophysical dust-gas mixtures. This scheme used a variable SPH smoothing length, SPH terms for the conservative part of the equations derived from a Lagrangian, and a `double-hump' shaped kernel that was found to improve the accuracy of the drag terms in both linear and non linear drag regimes. However, alongside the numerous tests we performed, we faced two rather fundamental difficulties which were not related to SPH but inherent to the two fluid formalism. Firstly, if dust particles concentrate below the resolution length of the gas phase, they tend to become trapped there, since they no longer feel any differential forces from the gas (see \\citealt{Ayliffe2012} for a discussion of this issue in SPH simulations; and \\citealt{PF10} for the same issue in a grid-based context). Secondly, when handling the specific problem of strong drag regimes (corresponding to small grains), we found two limitations which lead to a prohibitive computational cost: 1) the drastically small time steps required for the numerical stability of explicit schemes or the complexity of the implicit schemes involved, and less trivially 2) a high spatial resolution required to resolve the differential velocity between the gas and the dust in order to simulate the correct physical dissipation rate. The latter occurs because even if the differential velocity between the fluids is damped after a few stopping times $t_{\\rm s}$, the gas pressure causes a small spatial dephasing between the gas and dust. When the resolution is too low ($\\Delta x \\lesssim t_{\\rm s} c_{\\rm s}$, see \\citetalias{LP12a}), the dephasing from numerical simulations is artificially too large and the energy is over-dissipated.  This means that it is essentially impossible to simulate small grains accurately using the two fluid approach with standard fluids codes, since both infinitely small time steps and an infinitely large spatial resolution are required in the limit $t_{\\rm s}\\to 0$. Worse still, this limit corresponds to the rather obvious limit in which the two fluids are perfectly coupled and move precisely as a single fluid, albeit with a sound speed modified by the dust to gas ratio. For astrophysics this means that it is not currently possible to simulate the small, micron to cm sized grains accurately with \\emph{any} existing dust-gas code where dust is simulated using particles, and certainly not possible to simulate both small and large (metre-to-planetesimal sized) grains with the same technique.  In this paper, we show how the equations describing gas-dust mixtures can be reformulated to represent a single fluid moving with the barycentre of the mixture, leading to a set of equations only slightly modified from the usual equations of gas dynamics, with additional evolution equations for the differential velocity and the dust-to-gas ratio. This approach, though initially developed with the large drag/small grain regime in mind, turns out to be both general and elegant, since the important physical quantities of the mixture are computed directly, avoiding all of the artificial complications which arise in the two-fluid treatment.  The equations for the evolution of the single fluid are derived in Sec.~\\ref{sec:single} in primitive and conservative forms for the formalism to be relevant for both Lagrangian and Eulerian methods. This set of equations is completely general and can be used to simulate both large and small grains. In Sec.~\\ref{sec:approx}, we show how they can be further simplified in the specific limit of small grains and subsonic differential motion, leading to the standard equations of gas dynamics (with a modified sound speed), coupled with an advection-diffusion equation for the dust-to-gas ratio. In Sec.~\\ref{sec:tests} we demonstrate that the main physical effects associated with dust, including linear waves, shocks and the streaming instability, can all be captured with this simplified approach, and give the appropriate criterion for the use of the simplified formulation in numerical codes. ",
        "conclusions": "\\label{sec:conclu}  We have shown how the two fluid equations describing the evolution of a mixture of dust and gas can be reformulated in terms of a single fluid moving with the barycentric velocity of the mixture. The formulation consists of differential equations for the total mass $\\rho$, the barycentric velocity $\\vb$, the differential velocity $\\deltav$, and the dust to gas ratio $\\rhod/\\rhog$ (or equivalently, the dust fraction $\\rhod/\\rho$) that can be written in a form appropriate for both Lagrangian and Eulerian codes. The first two of these are identical (for $\\rho$) or only slight modifications (for $\\vb$) of the usual equations of gas dynamics. Evolving $\\deltav$ greatly simplifies the drag between the fluids, reducing it to a simple exponential decay, meaning that it is easy to solve this equation for both weak and strong drag regimes, and thus capture the dynamics of both small and large grains within the same formulation. Finally, explicit evolution of the dust to gas ratio means that it is straightforward to follow the concentration of solid material, which plays a crucial role in planet formation.  The one fluid approach solves with physics the two most fundamental issues related to numerical simulations of two fluid dust and gas mixtures. Firstly, the presence of only one resolution scale in the simulation means that the problem of over-concentration of one fluid below the resolution of the other cannot occur. Secondly, the equations reduce identically to single fluid gas dynamics in the limit of infinite drag, avoiding the need for both infinite spatial resolution and infinitesimally small timesteps that would be necessary with the two fluid approach.  We have also shown that strong drag regimes can be handled in an even simpler manner by adding a diffusion-like term in the equation governing the dust to gas ratio evolution. This approach was shown to capture most of the interesting physical processes in the mixtures, including the propagation of sound waves and shocks, and the linear growth of the streaming instability. This simplified formulation --- which can be implemented with only minor modifications to existing single fluid codes --- was shown to remain sufficiently accurate for use in numerical simulations provided the stopping time $\\ts$ is smaller than the minimum (Courant) timestep.  The nature of the evolution equations mean that this formalism lends itself readily to implementation in existing numerical codes. In a companion paper \\citep{LP14b}, we describe an implementation of the one-fluid formulation with the Smoothed Particle Hydrodynamics (SPH) method, though we stress that the approach is equally useful for both SPH and grid-based codes."
    },
    "1402/1402.3691_arXiv.txt": {
        "abstract": "It has at times been indicated that Landau introduced neutron stars in his classic paper of 1932. This is clearly impossible because the discovery of the neutron by Chadwick was submitted more than one month after Landau's work. Therefore, and according to his calculations, what Landau really did was to study white dwarfs, and the critical mass he obtained clearly matched the value derived by Stoner and later by Chandrasekhar. The birth of the concept of a neutron star is still today unclear. Clearly, in 1934, the work of Baade and Zwicky pointed to neutron stars as originating from supernovae. Oppenheimer in 1939 is also well known to have introduced general relativity (GR) in the study of neutron stars. The aim of this note is to point out that the crucial idea for treating the neutron star has been advanced in Newtonian theory by Gamow. However, this pioneering work was plagued by mistakes. The critical mass he should have obtained was $6.9\\,M_\\odot$, not the one he declared, namely, $1.5\\ M_\\odot$. Probably, he was taken to this result by the work of Landau on white dwarfs. We revise Gamow's calculation of the critical mass regarding calculational and conceptual aspects and discuss whether it is justified to consider it the first neutron-star critical mass. We compare Gamow's approach to other early and modern approaches to the problem. ",
        "introduction": "In Ref. \\cite{gamow} Gamow attempted to calculate the critical mass of a neutron star like object, approximating it with a homogenous matter distribution, as had previously been done by Stoner \\cite{stoner1} in the context of white dwarfs. Besides conceptual flaws, the derivation contains some calculational mistakes. The purpose of this work is to correct the latter and discuss the former. Gamow describes a state of matter at densities $\\rho \\approx 10^{12}\\mathrm{g/cm^3}$, where electrons have been absorbed by the $\\beta$-process and nuclear exchange forces come into play, which he ignores for the remainder of the calculation. He refers to this state as the {\\it nuclear state} and calls stellar cores in this state {\\it stellar nuclei}. The idea of Gamow's calculation is to compare the dependences of the Newtonian gravitational pressure \\begin{align} P_g=-\\frac{\\partial U_g}{\\partial V} \\end{align} and the non- and the ultra-relativistic degeneracy pressures \\begin{align} P_F=-\\frac{\\partial E_F}{\\partial V},\\ \\ P_{F,R}=-\\frac{\\partial E_{F,R}}{\\partial V} \\end{align} on the mass density $\\rho$. They are obtained for a homogeneous mass distribution, which appears to be a reasonable approximation for an order of magnitude estimate. Note that in this notation $P_F$ stands for the pressure of the Fermi gas and is not to be confused with the Fermi momentum. The total energies, including the rest-mass and the internal energies, but excluding the gravitational self-energy, are denoted $E_F$ and $E_{F,R}$, respectively, again not to be confused with the Fermi energy. In his calculation, Gamow obtains \\begin{align} P_g\\propto\\rho^{4/3},\\ P_F\\propto\\rho^{5/3},\\ P_{F,R}\\propto\\rho^{4/3}, \\end{align} and he concludes that for a non-relativistic Fermi gas, the degeneracy pressure will always balance the gravitational pressure as soon as a sufficiently high mass density is reached.  In the derivation of $P_F$, he assumes proportionality of mass density and number density, which is correct for a non-relativistic gas of Fermions. He also concludes that the fate of an ultra-relativistic mass distribution depends on the ratio $P_G/P_{F,R}$ at some initial time, as it will stay constant during a potential collapse. The flaw in his argument is the assumption of propotionality between the mass density and the number density in the derivation of $P_{F,R}$, which, for a free neutron gas, is not correct. In Section \\ref{sec:concFlaw}, we will argue that this assumption holds approximately true for a white dwarf with Fermi energy between the electron and deuterium mass regions, $m_e c^2\\ll E_\\mathrm{Ferm} \\ll 2 m_N c^2$. In the next section, we will turn to the calculational and conceptual flaws present in Ref. \\cite{gamow} and correct the numerical values obtained.  \\begin{table} \\begin{center} \\begin{tabular}{ccc} \\multicolumn{3}{c}{Homogenous models}\\\\ \\hline &$\\rho_\\mathrm{non}\\propto n$ & $\\rho_\\mathrm{ult}\\propto n^{4/3}$\\\\ \\hline\\hline $p_\\mathrm{non}\\propto n^{5/3}$\\ \\ \\  &\\ \\ no critical mass\\ \\ & not considered\\\\ $p_\\mathrm{ult}\\propto n^{4/3}$\\ \\ \\  & $6.9\\ M_\\odot$ & no critical mass\\\\ \\hline \\end{tabular} \\end{center} \\label{tab:homogen} \\caption{Critical masses in homogenous Newtonian models for all combinations of non- and ultra-relativistic approximations of $p$ and $\\rho$.} \\end{table}          \\begin{figure*} \\includegraphics[width=14.0cm]{./hydroNaive} \\label{fig:hydroNaive} \\caption{Relation between the total mass $M$ and the central density $\\rho_0$ for a non-relativistic mass density $\\rho$ and non-relativistic (solid), fully-relativistic (dashed) and ultra-relativistic (dotted) pressure $p$. The second model interpolates between the non- and the ultra-relativistic ones.} \\end{figure*} ",
        "conclusions": "The fact that numerical errors are present in Gamow's calculation alone would not suffice to consider it a failed attempt to calculate the neutron-star's critical mass, as it is an order of magnitude estimate, and does not take into account a radius-dependent density profile. However, because his approach is not applicable to the case of a free neutron gas, but rather to white dwarfs with ultra-relativistic electron pressures, and because he obtains the value of $1.5\\,M_\\odot$ by coincidence, while the corrected result reads $6.9\\,M_\\odot$, his result cannot be considered the first estimate of a neutron-star's critical mass. It should be considered a reproduction of earlier works \\cite{stoner1,anderson, stoner, chandrasekhar,chandrasekhar1,landau1} of Anderson, Stoner, Chandrasekhar, and Landau in the context of white dwarfs. The development of the first consistent neutron-star model should be attributed to Tolman, Oppenheimer, and Volkoff not only because it involves general relativitiy, but especially because it takes internal-energy contributions into account. Nevertheless, it was Gamow who initiated the transition of the neutron star from a theorized object \\cite{zwicky} to an object with predictable properties that could be tested by observation."
    },
    "1402/1402.1836_arXiv.txt": {
        "abstract": "We present the most up-to-date X-ray luminosity function (XLF) and absorption function of Active Galactic Nuclei (AGNs) over the redshift range from 0 to 5, utilizing the largest, highly complete sample ever available obtained from surveys performed with \\swift/BAT, \\maxi, \\asca, \\xmm, \\chandra, and \\rosat. The combined sample, including that of the \\subaru /\\xmm\\ Deep Survey, consists of 4039 detections in the soft (0.5--2 keV) and/or hard ($>2$ keV) band. We utilize a maximum likelihood method to reproduce the count-rate versus redshift distribution for each survey, by taking into account the evolution of the absorbed fraction, the contribution from Compton-thick (CTK) AGNs, and broad band spectra of AGNs including reflection components from tori based on the luminosity and redshift dependent unified scheme. We find that the shape of the XLF at $z \\sim 1-3$ is significantly different from that in the local universe, for which the luminosity dependent density evolution model gives much better description than the luminosity and density evolution model. These results establish the standard population synthesis model of the X-Ray Background (XRB), which well reproduces the source counts, the observed fractions of CTK AGNs, and the spectrum of the hard XRB. The number ratio of CTK AGNs to the absorbed Compton-thin (CTN) AGNs is constrained to be $\\approx$0.5--1.6 to produce the 20--50 keV XRB intensity within present uncertainties, by assuming that they follow the same evolution as CTN AGNs. The growth history of supermassive black holes is discussed based on the new AGN bolometric luminosity function. ",
        "introduction": "\\label{sec-intro}  Understanding the cosmological evolution of supermassive black hole (SMBHs) is a key issue in modern astrophysics. The good correlation of the mass of a SMBH in a galactic center with that of the bulge in the present-day universe \\citep[e.g.,][]{mag98,fer00,geb00,mar03,har04,hop07,kor09,gul09} indicates that SMBHs and galaxies co-evolved in the past. This idea is also supported by the similarity of the global history between star formation and SMBH growth \\citep[][]{mar04}. The so-called down-sizing or anti-hierarchical evolution, the trend that more massive systems formed in earlier cosmic time, has been revealed for both SMBHs (\\citealt{ued03}, U03; \\citealt{has05}, H05) and galaxies \\citep[e.g.,][]{cow96,kod04,fon09}.  Active galactic nuclei (AGNs) are the phenomena where SMBHs gain their masses from accreting gas by converting a part of their gravitational energies into radiation. It is known that the majority of AGNs are obscured by gas and dust surrounding the SMBHs, being classified as ``type-2'' AGNs. To elucidate the growth history of SMBHs, a complete survey of AGNs including heavily obscured populations throughout the history of the universe is necessary. X-ray observations, in particular those at high energies above a few keV, provide one of the most powerful approach for AGN detection thanks to the strong penetrating power against absorption and little contamination from star lights in the host galaxies. Furthermore, the deepest X-ray surveys currently available achieve the highest sensitivity even for unobscured (``type-1'') AGNs among those at any wavelengths \\citep[see][and references therein]{bra05}. The surface number density of the faintest X-ray AGNs reaches $\\sim 10^4$ deg$^{-2}$ \\citep{xue11}.  The integration of emission from all accreting SMBHs in the universe is observed as the X-ray background (XRB). To quantitatively solve the XRB origin is equivalent to revealing the cosmological evolution of AGNs that constitute the XRB. The XRB spans over a wide energy range from $\\sim$0.1 keV to $\\sim$100 keV and then is smoothly connected to the gamma-ray background at higher energies. Its spectrum has a peak energy density around $\\sim$20 keV. At energies below $\\sim$8 keV, now almost all of the XRB is resolved into discrete sources, mainly AGNs. Enormous efforts have been made to identify AGNs detected in X-ray surveys on the basis of multiwavelength observations, and the redshifts (and hence luminosities) of a large fraction of these sources are now estimated. These results make it possible to determine the spatial number density of AGNs that constitute the XRB below $\\sim$8 keV and its evolution. By contrast, in the hard X-ray band above $\\sim$10 keV, a significant fraction of the XRB is still left unresolved. Therefore, at present, the whole origin of the XRB over the wide range cannot be directly revealed by resolving individual sources.  It is very important to construct a ``population synthesis model'' of the XRB where the evolutions of all X-ray emitting AGNs with various types are formulated \\citep[for previous works, see e.g.][]{com95,gil99,ued03,bal06,gil07,tre09}. The model, in principle, must explain all the observational constraints including source number counts, redshift and luminosity distribution, the shape of the XRB. Once established, it gives the basis to understand the accretion history of the universe traced by X-rays, which is subject to least biases.  Two major elements in the population synthesis models are the X-ray luminosity function (XLF) and absorption function (or \\nh\\ function, \\citealt{ued03}) of AGNs. The XLF represents the number density of AGN per comoving space as a function of luminosity and redshift, one of the most important statistical quantity that can be determined from unbiased surveys.  Previously many studies have been made on the cosmological evolutions of the XLF of AGNs in the soft band below 2 keV \\citep[e.g.,][]{mac91,boy93,jon97,pag97,miy00,has05} and the hard band above 2 keV \\citep[e.g.,][]{ued03, laf05,bar05,sil08,ebr09,yen09,air10}. By using hard X-ray selected samples that contain both type-1 (unabsorbed) and type-2 (absorbed) AGNs, the absorption functions have been also investigated \\citep{ued03,laf05,bal06,tre06,has08}. Besides the well established anti-correlation of absorption fraction with luminosity \\citep[e.g.,][]{ued03,ste03,sim05}, several works have reported that the fraction of absorbed AGNs increased toward higher redshift from $z=0$ to $z > 1$ \\citep{laf05,bal06,tre06,has08}. More recent studies of high redshift AGNs at $z>2$ in deep fields (\\citealt{hir12}; \\citealt{iwa12}) reveal larger absorption fractions of high luminosity AGNs compared with the local universe. This redshift evolution is not included in the population synthesis model by \\citet{gil07}, one of the most widely referred model available at present.  There still remain uncertainties on the evolution of AGNs, however. The first issue is the shape of the XLF and its cosmological evolution. On the basis of hard X-ray surveys, \\citet{ued03}, \\citet{sil08}, \\citet{ebr09}, and \\citet{yen09} find that the XLF of AGN is best described with the luminosity dependent density evolution (LDDE) model. Later, \\citet{air10} propose that the luminosity and density evolution (LADE) where the shape of the XLF is constant over the whole redshift range unlike the case of LDDE also gives a similarly good fit to their data. While the down-sizing behavior is commonly seen in both models, quite different number of AGNs are predicted particularly at high redshifts of $z \\geq 3$. The second one is the number density (or fraction) of heavily obscured AGNs with \\nh\\ $>10^{24}$ cm$^{-2}$, so-called ``Compton-thick'' AGNs (CTK AGNs). Even hard X-ray surveys above 10 keV are subject to bias against detecting heavily CTK AGNs, because the transmitted emission is significantly suppressed due to repeated Compton scattering \\citep[see e.g.,][]{wil99,ike09}. It is not easy to identify individual CTK AGNs with limited photon statistics in deep survey data, and to estimate their intrinsic number density by correcting for such biases.  In this work, we present our latest results of AGN XLF over the redshift range from 0 to 5, by utilizing one of the largest combined sample ever available, obtained from surveys of various depth, width, and energy bands performed with \\swift /BAT, the Monitor of All-sky X-ray Image (\\maxi, \\citealt{mat09}), \\asca, \\xmm, \\chandra, and \\rosat. The sample consists of 4039 detections in the soft (0.5-2 keV) and/or hard ($>2$ keV) bands. We utilize a maximum likelihood (ML) method to reproduce the count-rate versus redshift distribution for each survey, by taking into account the selection biases. In the analysis, the contribution of CTK AGNs is considered, which is found to be more important in harder bands and at fainter fluxes. This enables us to determine the intrinsic XLF and the absorption function of type-1 plus type-2 AGNs with an unprecedented accuracy, thus establishing a standard population synthesis model of the XRB that is consistent with most of observational constraints currently available.  The organization of the paper is as follows. Section~\\ref{sec-sample} explains the sample used in our analysis. In Section~\\ref{sec-abs} the absorption properties of AGNs are discussed and the absorption function is formulated. Section~\\ref{sec-template} introduces the ``template model'' of broad band X-ray spectra of AGNs adopted in this work.  The distribution of photon index is examined by using the \\swift/BAT sample in Section~\\ref{sec-idx}. Section~\\ref{sec-lf} describes the details of main analysis using the whole sample. The best-fit results of the XLF are presented there. The predictions from the population synthesis model are given in Section~\\ref{sec-model}. We also discuss the constraints on the CTK AGN fraction and degeneracy with other parameters. Section~\\ref{sec-blf} represents an determination of the bolometric luminosity function of AGNs based on our new XLF, and the growth history of SMBHs. The conclusions of our work are summarized in Section~\\ref{sec-conclusions}. The cosmological parameters of \\cosmo\\ = (70$h_{70}$ km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7) are adopted throughout the paper. The ``log'' symbol represents the base-10 logarithm, while ``ln'' the natural logarithm. ",
        "conclusions": "\\label{sec-conclusions}  1. We have compiled so far the largest, highly complete sample of active galactic nuclei (AGNs) detected in X-ray surveys performed with \\swift/BAT, \\maxi, \\asca, \\xmm, \\chandra, and \\rosat, consisting of 4039 detections in the soft (0.5--2 keV) and/or hard ($>$ 2 keV) band. This gives us the best opportunity to trace the cosmological evolution of absorption properties and X-ray luminosity function (XLF) of all AGNs with log \\lx\\ (2--10 keV) = 42--46, including both type-1 (unabsorbed) and type-2 (absorbed) ones, from $z=0$ to $z=5$.  2. Using the latest compilation of spectral analysis of individual AGNs detected in the \\swift/BAT survey, we determine the shape of the absorption (\\nh) function in the local universe. We find that the fraction of absorbed AGNs with log \\nh\\ = 22--24 among all Compton-thin (CTN) AGNs with log \\nh\\ $\\leq$ 24 is 0.43$\\pm$0.03 at log \\lx\\ = 43.75. The distribution of photon index is peaked at $\\overline{\\Gamma_1}=1.94\\pm0.03$ for type-1 AGNs and $\\overline{\\Gamma_2}=1.84\\pm0.04$ for type-2 AGNs.  3. We confirm that the absorbed fraction of AGNs increases toward higher redshifts by keeping the anti-correlation with the luminosity. At log \\lx\\ = 43.75, the fraction of AGNs with log \\nh\\ = 22--24 among those with log \\nh\\ $\\leq$ 24 is proportional to $(1+z)^{0.48\\pm0.05}$ up to $z=2$.  4. To constrain the XLF of AGNs, we have developed a novel analysis method where we perform a maximum likelihood fit directly to the list of the observed count rate and redshift by taking into account selection biases in each survey. Here we consider the evolution of the absorbed fraction, the contribution of Compton-thick (CTK) AGNs with log \\nh\\ 24--26, and AGN broad band X-ray spectra including reflection components from tori based on the luminosity and redshift dependent unified scheme.  5. We find that the shape of the XLF at $z \\sim 1-3$ is significantly different from that in the local universe, showing flatter slopes in the lower luminosity range below the break. Its cosmological evolution is well described with the luminosity dependent density evolution (LDDE) model, while the luminosity and density evolution (LADE) model fails to fit the data.  6. On the basis of the absorption function and XLF determined above, we have newly constructed a population synthesis model of the X-Ray Background (XRB), which can be regarded as a ``standard model'' that well reproduces the source counts in both soft and hard bands, the observed fractions of Compton-thick AGNs, and the spectrum of the hard XRB.  7. To reproduce the hard XRB intensity in the 20--50 keV band within current uncertainties, we constrain that the fraction of CTK AGNs with log \\nh\\ = 24--26 to absorbed CTN AGNs with log \\nh\\ = 22--24 should be 0.5--1.6. This is also well consistent with the results of hard X-ray surveys above $\\simgt 10$ keV currently available.  8. We determine a bolometric luminosity function of AGNs by considering the luminosity-dependent bolometric correction factor and its variation from the XLF. The luminosity density of the whole AGNs has a peak around $z\\sim2$, where AGNs with bolometric luminosities of log $L = 46-47$ make the largest contribution. On the basis of Soltan's argument, the most recent estimate of the local mass density of supermassive black holes (SMBHs) is reproduced by adopting an averaged AGN radiation efficiency of $\\approx$0.05, although its mass dependence is suggested from the comparison of the AGN relic mass function and observed ones at $z=0$ and $z=1$."
    },
    "1402/1402.6627_arXiv.txt": {
        "abstract": "{We study the relationship between the {\\it K-band} and the sub-millimetre (submm) emissions of nearby galaxies by computing the bivariate {\\it K-band}-submm luminosity function (BLF) of the {\\textit{Herschel}}~\\thanks{\\textit{Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}} Reference Survey (HRS), a volume-limited sample observed in submm with {\\textit{Herschel}}/SPIRE. We derive the BLF from the {\\it K-band} and submm cumulative distributions using a {{\\it copula}} method. Using the BLF allows us to derive the relationship between the luminosities on more solid statistical ground. The analysis shows that over the whole HRS sample, no statistically meaningful conclusion can be derived for any relationship between the {\\it K-band} and the submm  luminosity. However, a very tight relationship between these luminosities is highlighted, by restricting our analysis to late-type galaxies. The fuminosity function of late-type galaxies computed in the {\\it K-band} and in the submm are dependent and the dependence is caused by the link, between the stellar mass and the cold dust mass, which has been already observed. } ",
        "introduction": "One of the most used tools in astronomy for probing the distribution of luminous matter in the Universe over cosmic time is the luminosity function (LF), which is the probability distribution function (PDF) of galaxy luminosities, the relative number of galaxies of different luminosities in a representative volume of the Universe \\citep[see e.g.][for a recent review of the subject]{Johnston11}. To accurately constrain the LF, a complete sample is required, a sample of objects selected at given frequencies with well defined selection criteria for which biases are understood well and accounted for. Biases are usually introduced when selecting a sample, if the selection criteria favour a particular class or population of objects. By incompleteness we mean a lack of a uniform coverage in redshift, in flux or in any other physical quantity \\citep[][]{SmithR12}.  To determine the LF, one starts from a sample selection at a given wavelength band and counts the number of galaxies in that band. If galaxies are selected in a particular band, but the number counts are carried out at a different wavelength, one needs to control possible selection effects in both bands before deriving any meaningful properties of the LF and its underlying sample \\citep[i.e.][]{Bose11}. In this case one can construct the  bivariate LF (BLF), which explores the dependency of the LFs at different wavelengths on each other, and explore the bivariate properties of some physical parameters of the sources \\citep{Johnston11}.  In a wider more general context, a bivariate distribution is often used to investigate the relationships between variables of a given sample. Quite often the bivariate distribution is obtained by ad hoc methods or heuristically, but analytic models are then required to interpret the results \\citep{Cholo85,Chap+03,Schaf07,CrossDriver02,Bal+06,Driver+06}. Inconsistent results for the relationship between two quantities are often obtained from regression analyses only because of the {\\it a priori} assumptions on the dependency between variables \\citep{LaFranca+95}. The difficulty in assigning an {\\it a-priori} distribution to a variable often results in a very poor fit from which no statistically meaningful conclusion can be derived \\citep{Bal+06}.  However, constructing a BLF from experimental  data is not a trivial task. Mathematically, this problem implies reconstructing a bivariate PDF satisfying prescribed marginals (i.e. cumulative distribution functions). If the distribution is not bivariate Gaussian, this operation is not straightforward. In fact, an infinite number of bivariate distributions exists with the same marginals, which can be constructed once the dependence structure is specified. The goal is, therefore, to find a way to disentangle the marginal distribution and the dependence structure. Here we adopt an approach based on the PDF of each random variable and the so-called {\\rm copula}, a function that provides the dependence structures between them. We show in detail how it is possible to reconstruct a BLF from a multi-wavelength dataset.  In the past the BLF has been computed at optical/X-ray wavelengths by La Franca et al. (1995) % and for the SDSS sample by Cross \\& Driver (2002), Ball et al. (2006), and Driver et al. (2006). % All these approaches were based on a maximum likelihood (ML) fit to a bivariate PDF, or its non-parametric version (the stepwise ML fit SWML introduced by Efstathiou et al. 1988). % However, until recently there has not been any general method of constructing a bivariate distribution function with pre-defined marginal distributions and correlation coefficient. Our approach is fundamentally different from previous works dealing with the bivariate LF. We follow Takeuchi (2010), % who introduced a {\\rm copula} method to estimate the bivariate LF in the infrared, but we use a different implementation from these authors, mainly in the computation of the correlation coefficient.   We compute the bivariate {\\it K-band}-submm luminosity function of the {\\textit{Herschel}} Reference Survey (HRS) sample \\citep{Bose+10a}. This sample has been extracted from the 2MASS survey and observed in the submm by the SPIRE instrument on board of the {\\textit{Herschel}} satellite. We make use of an analytical function of the {\\it K-band} and submm LFs, and we use it as the marginal PDF. As explained in the text, the choice we made is dictated by the smallness of the sample and by the presence of upper limits among the SPIRE observations.  Since our sample is inhomogeneous and relatively small, the direct derivation of the submm LF cannot be pursued \\citep[as in e.g.][]{Dye+10, Dunne+11, Davies+10}, while we can exploit the knowledge of the {\\it K-band} LF, which on the contrary is well established \\citep{Cole+01,Kochanek+01}. In addition using the BLF as a statistical tool in the presence of upper limits provides results that lie on more solid statistical ground than any other simpler tool, i.e. a linear regression test. %  The paper is organised as follows. The methodology for computing the luminosity functions is described in section~\\ref{tools}. The sample selection is briefly described in section~\\ref{sample} with the results shown in section~\\ref{results} and a discussion in section ~\\ref{discussion}. ",
        "conclusions": "The relationship between the {\\it K-band} and the far infrared (submm) emissions of nearby galaxies, investigated through the BLF [{\\it K-band}, submm] of the {\\textit{Herschel}} Reference Survey (HRS), shows that while a poor correlation is found using a whole HRS sample. The late-type galaxies subsample presents a remarkable correlation, which reflects on the physical relationship between the two frequency ranges. These results suggest that the LF of galaxies computed in the {\\it K-band} and in the submm are dependent and the dependence lies in the relationship between the galaxy stellar mass and the cold dust mass.  Finally, the agreement for the estimate of the LF of the HRS sample with what was estimated for the parent sample (2MASS survey) further supports the conclusion that inhomogeneities associated with large scale structure in the local Universe does not affect the  statistical properties of the galaxy distribution much. This conclusion, together with our new statistical analysis, has allowed us to better characterise the HRS sample, where, despite the in-homogenous selection and a partial bias towards low submm luminosities, our sample can reasonably be considered as representative of a nearby population of galaxies."
    },
    "1402/1402.0600_arXiv.txt": {
        "abstract": "Self-gravitating Bose-Einstein condensates have been proposed in various astrophysical contexts, including Bose-stars and BEC dark matter halos. These systems are described by a combination of the Gross-Pitaevskii and Poisson equations (the GPP system). In the analysis of these hypothetical objects, the Thomas-Fermi (TF) approximation is widely used. This approximation is based on the assumption that in the presence of a large number of particles, the kinetic term in the Gross-Pitaevskii energy functional can be neglected, yet this assumption is violated near the condensate surface. We also show that the total energy of the self-gravitating condensate in the TF-approximation is positive.  The stability of a self-gravitating system is dependent on the total energy being negative. Therefore, the TF approximation is ill suited to formulate initial conditions in numerical simulations. As an alternative, we offer an approximate solution of the full GPP system. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.2096_arXiv.txt": {
        "abstract": " ",
        "introduction": "The year 2012 marks the hundredth anniversary of the detection of cosmic rays by Viktor Hess. As we know today, cosmic rays consist essentially of protons and nuclei of heavier elements; electrons contribute only at the percent level.  Since cosmic rays are electrically charged, they are deflected by cosmic magnetic fields on their way to Earth. Precise pointing -- i.e. astronomy -- is only possible with electrically neutral, stable particles: electromagnetic waves (i.e. gamma rays at the energies under consideration) and neutrinos. High energy neutrinos, with energies much beyond a GeV, must be emitted as a by-product of collisions of charged cosmic rays with matter. Actually, only neutrinos provide incontrovertible evidence for acceleration of hadrons since gamma rays may also stem from inverse Compton scattering of accelerated electrons and other electromagnetic processes. Moreover, neutrinos can escape much denser celestial environments than light, therefore they can be tracers of processes which stay hidden to traditional and gamma ray astronomy. At the same time, however, their extremely low reaction cross section  makes their detection a challenge and requires huge detectors.  This article follows the evolution towards IceCube, the first neutrino telescope with a realistic discovery potential, sketches the recent relevant results obtained with large neutrino telescopes and tries a look into the future. Recent reviews of the field can be found in \\cite{reviews,KS}, a detailed review of its history in \\cite{history}. ",
        "conclusions": ""
    },
    "1402/1402.3745_arXiv.txt": {
        "abstract": "The acceleration of charged particles is relevant to the solar corona over a broad range of scales and energies. High-energy particles are usually detected in concomitance with large energy release events like solar eruptions and flares, nevertheless acceleration can occur at smaller scales, characterized by dynamical activity near current sheets. To gain insight into the complex scenario of coronal charged particle acceleration, we investigate the properties of acceleration with a test-particle approach using three-dimensional magnetohydrodynamic (MHD) models. These are obtained from direct solutions of the reduced MHD equations, well suited for a plasma embedded in a strong axial magnetic field, relevant to the inner heliosphere. A multi-box, multi-scale technique is used to solve the equations of motion for protons. This method allows us to resolve an extended range of scales present in the system, namely from the ion inertial scale of the order of a meter up to macroscopic scales of the order of $10$\\,km ($1/100$th of the outer scale of the system). This new technique is useful to identify the mechanisms that, acting at different scales, are responsible for acceleration to high energies of a small fraction of the particles in the coronal plasma. We report results that describe acceleration at different stages over a broad range of time, length and energy scales. ",
        "introduction": "Understanding the mechanisms that accelerate charged particles in dynamical plasmas is an intense subject of research, with applications to astrophysical, space, and laboratory environments \\citep{cvb12}. Charged particles are energized in solar flares, planetary magnetospheres, at interplanetary shocks, and in the interstellar medium. In the solar corona electrons can be accelerated to tens of MeV, while ions gain energies up to several GeV, and high-energy particles may carry up to $30 \\%$ of the energy budget of a flare \\cite[for a review see][]{asc02, km07}. The fast release of energy in solar flares is commonly attributed to magnetic reconnection events. In such complex natural systems, several processes may contribute to the observed particle energization, and these may be distinguished in part by the locations and scales at which they occur. In this way the resulting acceleration might be viewed as a \\textit{multi-stage process} \\cite[see, e.g.,][]{jag86, gor05, bro09}.  In particular the role of scattering from magnetic irregularities is well known to be an important factor \\citep{Jokipii66}: particles can be accelerated by the large scale {\\bf electric} fields \\citep{Speiser65, Litvinenko96}, as well as by the smaller scales of reconnecting fields \\citep{DrakeEA05} and/or discontinuities, including acceleration in shocks \\citep{Giacalone&Jokipii96, Giacalone05}. Furthermore they are also accelerated by resonances with structures at various scales, and acceleration mechanisms may be strongly influenced in various contexts by changes of magnetic connectivity \\citep[e.g.,][]{rmr12}. Thus essentially all length scales may be important in the investigation of the energization process, giving rise to a rather complex behavior in the case of turbulent fields, characterized by the presence of energy at all scales from large to small and of current and vorticity coherent structures, that can cause resonant or nonresonant interactions \\citep{MatthaeusEA84, KobakEA00}. Indeed, although the smooth or average electromagnetic fields are associated with a variety of possible particle motions \\citep{Speiser65, Sonnerup71, Cowley78}, magnetic fluctuations can strongly influence both the local electric field and particle motion. For example strong resonant pitch angle scattering, as well as the formation of transient trapping centers, such as secondary islands, are directly associated with the presence of electromagnetic fluctuations \\citep{MatthaeusLamkin86, RuffoloEA03, DrakeEA06}.   Such effects are expected to be widespread since turbulent magnetic fluctuations are observed in space plasmas in practically  all environments. Fully developed turbulence involves itself a hierarchical process in which many scales of motion are engaged in the energy transfer. As particles become energized, their interactions with the turbulence sample a correspondingly wide range of scales. From an observational point of view, different spectra and intensity variations have been detected from sources of nonthermal radiation in many flares \\cite[see, e.g.,][]{tak95, ems03, fle07}, which is also hard to explain invoking only a single acceleration mechanism.  Because of the complexity of particle motion, a useful approach has been to resort to numerical experiments based on test particle simulations, in which no interaction among the particles and no back reaction to the imposed electromagnetic fields is considered \\citep{LitvinenkoEA93, TurkmaniEA05, DallaEA05, OnofriEA06}. However, processes that span a wide range of scales pose profound challenges because the different scales require different time and spatial resolutions, and qualitatively different phenomenologies may apply. For instance, in the description of a coronal loop the typical scale of the system $L$ measures thousands of kilometers. In contrast the ion inertial scale $d_i$, the typical scale at which dissipative effects become important \\citep{SmithEA01}, is of the order of a meter and is comparable with the typical gyroradius, $\\rho = v_\\textrm{th}/\\Omega$, of a particle moving at the thermal speed $v_\\textrm{th}$ with a gyrofrequency $\\Omega$.  Such conditions, spanning a wide range of length scales, are typical in space plasmas. For example, in the solar wind it can be estimated that $L/d_i \\sim 10^5$. In a coronal hole (at about one solar radius) and in magnetically confined regions of the solar corona $L/d_i \\sim 10^6$. Attempting to resolve this full range of scales requires an enormous cost in terms of computing power at present standards. In galactic astrophysics even larger $L/d_i$ may be relevant. To attack such problems, approximate models, such as the guiding center and gyrokinetics, have been employed to describe particle motion in a strong magnetic field \\cite[see, e.g.,][]{lee83, bv01, giu05, gor05, LeheEA09}. These models are derived using magnetic moment $\\mu$ conservation, so their validity is limited to specific circumstances. In fact the equation of motion of the particle guiding center can be solved in the case of small gyroradius and large gyrofrequency, $\\rho \\ll l$ and $\\Omega \\gg T^{-1}$, where $l$ and $T$ should be interpreted as the typical length and time scales of the {\\it local} electromagnetic field. As a consequence, even if particles are strongly magnetized globally, non-adiabatic behavior may develop at local scales \\citep{bro09, dal12a}, in particular in the presence of sharp field gradients, near reconnection sites or in  presence of well developed magnetohydrodynamic turbulence. In these situations the presence of resonant effects and/or coherent structures could violate magnetic moment $\\mu$ conservation. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{fig1.eps} \\end{center} \\caption{Schematic representation of the hierarchical computational boxes used in the multi-scale model. The boxes are progressively larger to accommodate higher energy particles (with correspondingly larger giroradii). Physical fields in a larger box are obtained rescaling those from the smaller box. Indicating, e.g.,  with $L$ and $\\Delta x$ respectively the linear length and grid resolution along the $x$-direction of the smaller box, the fields in the larger box of linear length $L'$ and grid size $\\Delta x' = \\Delta x\\, L'/L$ are built rescaling the fields from the smaller box by a factor $L'/L$ in each linear direction and remapping them to the bigger box.}  \\label{fig:box} \\end{figure}  To represent turbulent fields at all relevant scales, one approach is to use gridless methods with prescribed analytical fields \\cite[see, e.g.,][]{Giacalone&Jokipii96, qin02, byk08, fras08, fras12}. In principle, this method can represent a wide range of scales by correspondingly increasing the number of modes retained in the representation. However, the direct gridless implementation becomes extremely onerous in that case. These models also most often employ random phases, which prevent representation of coherent small scale structures, such as current sheets and vorticity concentrations, that can play a key role in particle acceleration.   Therefore we have developed a new multi-box, multi-scale test-particle technique to resolve the wide range of time and length scales usually involved in the description of space and astrophysical plasmas. The main idea is to study particles dynamics at different steps during their evolution, moving from several simulation boxes of different dimensions (for a schematic picture see Figure~\\ref{fig:box}). The nonrelativistic equations of motion of a particle moving in an electromagnetic field are solved numerically, keeping approximately constant the resolution of the particle {\\bf gyroradius} $\\rho$ going from a smaller to a larger box. This has been accomplished choosing for every box a grid scale $\\Delta x$ that satisfies the condition $\\Delta x < \\rho$. Different spatial and temporal resolutions in different regions allow us to treat with the same accuracy all the scales involved in the system. In this paper we present this new numerical technique applied to the study of particle acceleration in a coronal loop. We report results that describe acceleration in direct simulations of reduced magnetohydrodynamic (RMHD) turbulence at different stages, over a broad range of time, length and energy scales.   The paper is organized as follows. In Section~\\ref{sec:2} we describe the equations and properties of the MHD fields and of the particles. In Section~\\ref{sec:3} the multi-box technique is introduced. In Section~\\ref{sec:4} results are given for energization of protons obeying the nonrelativistic equations of motion. Section~\\ref{sec:5} contains our conclusions. ",
        "conclusions": "\\label{sec:5}  In this paper we investigated particle acceleration in turbulent electromagnetic fields embedded in a strong guide magnetic field, an important issue in astrophysical, space and laboratory plasmas. We performed test particle simulations for protons in turbulent fields obtained from direct numerical solutions of reduced Magnetohydrodynamic turbulence (for a detailed description of the RMHD simulation see \\cite{rved08}). Although the test particle approach is not self-consistent, it provides a useful bridge between the large-scale macroscopic MHD and a more realistic particle kinetic physics description. The geometry of the system and the typical parameters used in the simulations are well-suited for the description of particle acceleration in coronal loops.  A common problem in the numerical description of plasma and/or particle dynamics in space plasmas is the involvement of a huge range of length and time scales. In the case of coronal loops the large scale motion occurs on a distance comparable with the loop perpendicular length (thousands kilometers), but the acceleration process takes place at different stages and/or locations during the evolution of the system and may require the participation of several mechanisms. Therefore it can be considered as a multi-stage process, and the numerical study may require that the various length and time scales should be followed with different time and spatial resolutions. In the study of this kind of system, the acceleration may be associated with strong and sudden changes, including changes of magnetic topology near reconnection sites as well as strong discontinuities. Consequently the use of the guiding center equations may be inappropriate, because non-adiabatic behavior may develop on local scales.  Motivated by these reasons, and to maintain the possibility of gyro-resonant interactions at all phases of the acceleration process, we have developed a {\\it multi-box, multi-scale} technique to solve the nonrelativistic equation of motion of a particle moving by the action of electromagnetic fields (see Section~\\ref{sec:3}). This technique allowed us to resolve an extended range of scales present in the system, namely from the ion inertial scale $d_i$, of the order of $1$\\,m, up to macroscopic scales of the order of $10$\\,km ($1/100$th of the correlation scale, the photospheric convection scale $\\sim 10^3$\\,km), resolving at the same time the particle gyroradius. An important feature of the multi-box method is that the scaling of the MHD fields is controlled so that the magnetic field, plasma velocity, and electric field take on similar values in the rescaled boxes as their size varies from the initial smallest box to the largest final box. This amounts to changing length scales and resistivities in the sequence of boxes to maintain the required statistical continuity.   Employing the multi-box strategy, we performed a numerical experiment that shows how acceleration takes place typically in a two stage process. First, a small fraction of initially loaded low energy particles are entrained in elongated current sheets long enough to gain substantial energy in response to the parallel electric field in the current sheet. High parallel speeds are attained in this stage. Gradually these particles escape the current channels due to pitch angle scattering and due to their larger gyroradius after energization. In the second stage, particles encounter nonuniform in-plane electric fields that enable a resonant betatron acceleration for those particles that remain entrained in these regions for a suitable amount of time. In this case the acceleration is mainly perpendicular. However once again pitch angle scattering tend to isotropize the particles, which also then facilitates their escape from  the accelerating region.  A point of interest is that the two stage acceleration process described here is fully consistent with important features of test particle acceleration in turbulence that have been reported previously. For example, the initial parallel acceleration of test particles in current channels associated with two dimensional reconnection was reported some time ago by \\cite{amb88}. There, too, it was observed that the energization is limited by the time of temporary trapping of particles near the acceleration region. The role of pitch angle scattering in disrupting the parallel acceleration process was also described. Furthermore it is likely that secondary islands near or in current sheets may enhance the ability of sheet-like structures to entrain and accelerate particles \\citep{amb88,DrakeEA06}. We have not attempted to describe the effect of secondary islands here.  It is also noteworthy that the parallel and perpendicular stages of acceleration that we described were anticipated, although as separate effects, in the study by \\cite{dmt04}. In that case, using a single simulation box, it was shown that at early times, protons experience perpendicular acceleration, while electrons experience parallel acceleration. The consistency of that picture with the present one is established by noting that in that study, particles were initialized with zero speed, but quickly acquire speeds comparable to the Alfv\\'en speed. However the current sheets have thickness comparable to the ion inertial scale, so the protons almost immediately encounter the edges of the current sheets and therefore enter the perpendicular ``stage two'' process described here. Given their much lower inertia, electrons have very small gyroradii when moving at the Alfv\\'en speed, and so can experience what we presently call the stage one parallel acceleration process, gaining substantial energy in a short time. The multi-box method allows the computation to follow both stages one and two for a single particle type, while maintaining resonant couplings that are central to spatial transport and therefore contributing to entrainment and escape processes.  For the particular example appropriate to coronal loops that we showed here, by the end of the third box in the multi-box sequence particles attained energies up to $0.1$ Gev. In order to cover a realistic large parallel coronal loop length, another two simulations in larger boxes would need to be performed. However, at the end of the third simulations particles have reached energies of 0.1 GeV and a non-relativistic approach could not be not suitable for higher energy. Continuing in this way, higher energies are expected at the end of the largest scale simulation, and  an ``escape'' mechanism mimicking the collision of particles with the denser chromospheric layer \\citep[e.g.,][]{GontiEA13} should be included. Furthermore it must be pointed out that the lack of self-consistency intrinsic in test-particle simulations could become at some point an important limitation too, especially for the highest energy particles.  Although a refinement of the {\\it multi-box} technique may be also required, the results shown in this paper indicate a promising way of understanding energetic particles in astrophysical, space, and laboratory environments and studying some basic issues in the complex topic of bridging the MHD and kinetic descriptions of a plasma."
    },
    "1402/1402.1576_arXiv.txt": {
        "abstract": "We present fully non-perturbative quantum Monte Carlo calculations with non-local chiral effective field theory (EFT) interactions for the ground state properties of neutron matter. The equation of state, the nucleon chemical potentials and the momentum distribution in pure neutron matter up to one and a half times the nuclear saturation density are computed with a newly optimized chiral EFT interaction at next-to-next-to-leading order. This work opens the way to systematic order by order benchmarking of chiral EFT interactions, and \\emph{ab initio} prediction of nuclear properties while respecting the symmetries of quantum chromodynamics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.4509_arXiv.txt": {
        "abstract": "We report the detection of interstellar silicate dust in the z$_{abs}$=0.685 absorber along the sightline toward the gravitationally lensed blazar TXS 0218+357. Using \\textit{Spitzer Space Telescope} \\textit{Infrared Spectrograph} data we detect the 10~$\\micron$ silicate absorption feature with a detection significance of 10.7$\\sigma$. We fit laboratory-derived silicate dust profile templates obtained from literature to the observed 10~$\\micron$ absorption feature, and find that the best single-mineral fit is obtained using an amorphous olivine template with a measured peak optical depth of $\\tau_{10}$=0.49$\\pm$0.02, which rises to $\\tau_{10}\\sim$0.67$\\pm$0.04 if the covering factor is taken into account. We also detected the 18~$\\micron$ silicate absorption feature in our data with a $>$3$\\sigma$ significance. Due to the proximity of the 18~$\\micron$ absorption feature to the edge of our covered spectral range, and associated uncertainty about the shape of the quasar continuum normalization near 18~$\\micron$, we do not independently fit this feature. We find, however, that the shape and depth of the 18~$\\micron$~silicate absorption are well-matched to the amorphous olivine template prediction, given the optical depth inferred for the 10~$\\micron$ feature. The measured 10~$\\micron$ peak optical depth in this absorber is significantly higher than those found in previously studied quasar absorption systems. The reddening, 21-cm absorption, and velocity spread of Mg~II are not outliers relative to other studied absorption systems, however. This high optical depth may be evidence for variations in dust grain properties in the ISM between this and the previously studied high redshift galaxies. ",
        "introduction": "The dust and gas that comprise the interstellar medium (ISM) play an important role in galaxy evolution, despite the fact that they constitute a relatively small fraction of each galaxy's total mass (e.g., $\\sim$10\\% in the Milky Way). The presence of interstellar dust is physically manifested both through absorption and scattering processes which redistribute $>$30\\% of the starlight emitted at UV-optical wavelengths over a range of infrared (IR) wavelengths \\citep{Bernstein02}, and through the depletion of refractory elements from ISM gas. The physical and chemical states of the ISM are significantly affected by the dust, which regulates the thermal energy balance through the ejection of photoelectrons in photodissociation regions, and via IR emission from molecular and shock-heated plasmas. Dust grain surfaces also promote the formation of molecules, such as H$_2$, which act as a coolant facilitating molecular cloud collapse. As the ISM gas is formed into new stars and the galaxy's stellar population evolves, the ratio of the silicate dust to the carbonaceous dust reflects the relative abundances of their different stellar sources. Silicate dust is typically produced in oxygen-rich environments, such as near Type II supernovae or AGB stars with low C/O ratios, while carbonaceous dust is associated with carbon-rich environments, such as near evolved low-mass stars, although the details of the formation are still debated both in local and in higher redshift galaxies [e.g., \\citet{Dwek11}]. Since dust is relatively quickly destroyed through shock processing in the ISM, with a typical lifetime of $\\approx$500~Myr for many dust species \\citep{Jones96,Jones11}, interstellar dust is a sensitive probe of the state of the ISM material and stellar populations at a given epoch, assuming that selective destruction and reformation rates in the ISM are not highly variable.  Quasar absorption systems (QASs) provide a unique tool to simultaneously probe the dust and gas content in the ISM of distant galaxies, independent of galaxies' intrinsic luminosities or star-formation rates. The study of QASs exploits sightlines to distant, luminous objects, such as quasars, that emit radiation over a wide-range of frequencies. As the quasar light passes through foreground galaxies, interstellar gas and dust produce absorption features in the quasar spectrum. The ISM of any galaxy with a luminous background source can be studied with this technique, but it is particularly effective when applied to the most gas-rich systems, such as damped Lyman-$\\alpha$ absorbers (DLAs; neutral hydrogen column densities N$_{HI}\\geq2\\times10^{20}$ cm$^{-2}$) and sub-DLAs (10$^{19}\\leq$N$_{HI}<2\\times10^{20}$ cm$^{-2}$). These gas-rich QASs are the primary neutral gas reservoir for star formation in the universe [e.g., \\citet{Storrie00, Peroux05, Prochaska05}]. As opposed to flux-limited studies of gas and dust in emission, that are biased toward the most robustly star-forming objects, QASs provide insight into more quiescent or moderately-star-forming galaxies, potentially more similar to Local Group galaxies.  Evidence for the presence of interstellar dust in gas-rich QASs has been found in both moderate-to-large surveys, as well as in investigations of individual systems. Studies of QASs show depletions of refractory elements and reddening of background quasars \\citep{Pei91, Pettini97}, as well as a correlation between dust-reddening and gas metal absorption line strengths in absorbers at 1$<$z$<$2 \\citep{York06}. The bulk of analyses exploring dust spectral properties in QASs, however, concentrate on signatures associated with the carbonaceous dust component, such as the 2175~\\AA~bump, rather than on the silicate dust, which exhibits prominent IR features near 10 and 18~$\\micron$ [see, e.g., \\citet{Draine03} for review of dust spectral signatures and their origins]. It is important to also consider this silicate dust, both because it is primarily produced by a different stellar sub-population, and because within the Milky Way $\\sim$66\\% of the core mass of interstellar dust grains is silicate in composition \\citep{Greenberg99, Zubko04}.  In an exploratory study, our group made the first measurement of silicate dust in a QAS. A $>$10-$\\sigma$ detection of the 9.7~$\\mu$m silicate absorption feature was made in the z$_{abs}=0.524$ DLA toward AO 0235+164 using low-resolution \\textit{Spitzer Space Telescope Infrared Spectrograph (IRS)} data. This was followed by identification of the 9.7~$\\micron$~silicate absorption feature in 5 additional strong QASs at redshifts 0.44$<z_{abs}<$1.31 \\citep{Kulkarni11, Aller12}. Analyses of the silicate dust in these 6 QASs suggest both similarities and differences relative to that in the Milky Way ISM. In the Milky Way diffuse interstellar clouds, the peak optical depth of the silicate absorption is seen to scale with the extinction ($A_V$), and hence with the reddening [E(B-V)=$A_V/R_V$ where $R_V$, the selective extinction, is $\\sim$3.1]. In the studied QASs, the silicate optical depth similarly scales with the reddening. However, the absorption features are 2-4 times deeper for a given reddening than those in the Milky Way \\citep{Kulkarni11}. This disparity may stem from differences in dust composition, metallicity, and/or grain size distributions. A larger selective extinction may imply, for instance, larger dust grain sizes. Furthermore, in these 6 QASs comparisons of the peak wavelengths of the 10~$\\micron$~silicate absorption feature, and of the substructure within the feature, reveal variations suggestive of mineralogical and grain structural differences relative to one-another, and relative to the Milky Way ISM \\citep{Kulkarni11,Aller12}. These differences are most pronounced in the z=0.886 absorber toward PKS 1830-211 \\citep{Aller12}, where the 10~$\\micron$ absorption feature substructure is consistent with a crystalline olivine silicate dust grain origin, in marked contrast with the amorphous ($<$5\\% crystalline) silicate dust grains in the Milky Way ISM \\citep{Kemper04, Li07}. Exploring whether these trends and variations are found in other QASs is essential for quantifying the evolution in dust grain properties.  In this paper, we make progress toward this goal by presenting the first analysis of the 10 and 18~$\\micron$ silicate absorption features in the z$_{abs}$=0.68466$\\pm$0.00004 absorber \\citep{Carilli93} toward the gravitationally lensed, z=0.944$\\pm$0.002 blazar TXS 0218+357 at RA (J2000)= 02\\fh21\\fm05\\fs5~and Dec (J2000) = +35\\degr56\\arcmin14\\arcsec~ \\citep{Patnaik92,Browne93,Cohen03}. Throughout our analysis we simply refer to this z$_{abs}$=0.6847 system, which is by far the strongest known metal absorption system along the quasar sightline, as \\textit{the} TXS 0218+357 QAS. The TXS 0218+357 QAS was selected for this study for several reasons.  First, this QAS is relatively rich in both gas and dust. The abundance of gas is evident from the large measured Mg~II absorption rest-frame equivalent width W$_{2796}=$3.0~\\AA, the high molecular hydrogen column density ($5\\times10^{21} \\leq N_{H2}\\leq 5\\times10^{23}$ cm$^{-2}$), and the detections of 21-cm absorption implying $N(H I) = 4\\times10^{18} (T_s/f)$ cm$^{-2}$, where $T_s$ is the gas spin temperature and $f$ is the H~I covering factor \\citep{Browne93,Kanekar02,Combes95,Menten96,Carilli93}. There are also signatures of dust in the QAS, including an exceedingly high measured reddening of the background quasar. Statistical studies of mean DLA extinction curves using SDSS data find that most QASs at 1$<$z$<$5 have a typical rest-frame reddening of E(B-V)$\\sim$0.002 \\citep{York06,Khare12}, while the TXS 0218+357 QAS exhibits a total measured E(B-V)=0.62$\\pm$0.04, assuming R$_V=3.1$ \\citep{Falco99}.  Second, this QAS is known to be rich in molecules, similar to the PKS 1830-211 z=0.886 absorber. Since in the PKS 1830-211 QAS the interstellar silicate dust may be significantly crystalline \\citep{Aller12}, unlike that in the Milky Way, it is important to ascertain whether other molecule-rich QASs are likewise potentially rich in crystalline silicates. Molecules such as CO, CH, and H$_2$ trace cold, dense, star-forming material in the galaxy ISM. The TXS 0218+357 QAS exhibits transitions for numerous molecules including CO, HCO$^+$, HCN, HNC, CN, CS, C$^{18}$O, $^{13}$CO, H$_2$O, NH$_3$, OH, H$_2$CO and its anti inversion, and the primordial LiH molecule which traces H$_2$ in low metallicity regions \\citep{Wiklind95, Combes95, Combes97b, Combes98,Combes97,Gerin97,Menten96,Kanekar02,Kanekar03,Henkel05,Jethava07,Zeiger10,Friedel11}. Moreover, this absorber has among the highest molecular column densities at this redshift range \\citep{Combes98}. Studies of the QAS reveal that one of the gravitationally-lensed quasar sightlines passes through an optically thick and clumpy molecular cloud, while the other sightline passes through a more optically thin region \\citep{Combes95,Combes97b}. The derived column densities and abundances for molecules such as formaldehyde along the more obscured sightline are similar to those found in Milky Way molecular clouds \\citep{Menten96}, while the NH$_3$ absorption measurements are more consistent with a diffuse cloud origin \\citep{Henkel05}.  Third, the z$_{abs}$=0.68466 QAS redshift  is optimal for studying both the 10 and 18~$\\micron$ silicate absorption features with the \\textit{IRS} instrument used in our previous studies. In the QAS rest-frame, the IRS spectra span the region 3.1-22.6~$\\micron$, with adequate quasar continuum coverage between the absorption features. Additionally, for this quasar the \\textit{IRS} sensitivity was sufficient to obtain data for this analysis in only 30 minutes of exposure time.  Fourth, the z$_{em}$=0.944 \\citep{Cohen03} background quasar is gravitationally lensed, which boosts its signal, and hence the signal-to-noise ratio (S/N) of the QAS gas and silicate dust absorption features. Furthermore, since this is a well-studied strong lens system, information about the QAS host galaxy topography, which is often lacking for other QASs, is available. Studies of the morphology of the lensing galaxy hosting the QAS indicate it is a gas-rich spiral galaxy with narrow absorption (Ca~II H\\&K, Mg~II, Fe~II, Mg~I) and emission ([OII], H$\\beta$, [OIII]) lines \\citep{Carilli93,Browne93,Menten96,York05}, suggestive of some ongoing star-formation, comparable to that in spiral galaxies like the Milky Way based on the measured [OII] equivalent width \\citep{Browne93}. Radio studies [e.g., \\citet{Odea92,Patnaik93}] reveal that the quasar is being lensed into two point sources and a small Einstein ring with an angular diameter of 335 mas ($\\sim$2.4~kpc in the lensing galaxy for H$_o$=70 km/s/Mpc). The optical and IR magnitudes along the two compact-object lines of sight imply a reddened and optically faint (V$>$21mag) source \\citep{Lehar00}.  Lastly, owing to the well-studied strong gravitational lens, as well as to the monitored (variable) blazar, abundant ancillary data are available for this system. These data include spectral energy distributions (SEDs) extending from the $\\gamma$ ray (e.g. Fermi, ROSAT, Swift) through the optical and NIR (e.g. HST), to the sub-mm (e.g. NRAO, PdBI) and radio (e.g., VLA, VLBA). However, while these data provide global constraints, owing to the blazar variability, predicting the flux at a given epoch is challenging.  In this paper, we begin with a description of the \\textit{IRS} observations, followed by a summary of the data reduction (\\S\\ref{IRSspect}), spectral extractions (\\S\\ref{Extract}), and quasar continuum normalizations (\\S\\ref{Norm}). We then discuss the profile template fitting to the silicate absorption features produced by the TXS 0218+357 QAS (\\S\\ref{results}). We conclude with a discussion of our findings in comparison with the silicate dust in other QASs (\\S\\ref{disc}), followed by a summary of our main results (\\S\\ref{sum}). In the Appendices, we show the effects of different quasar continuum normalizations on our results (Appendix~\\ref{normAP}), and discuss details of the template profiles applied to the data (Appendix~\\ref{templateAP}). ",
        "conclusions": "\\label{disc} We find that the best-single mineral fit to the silicate absorption features in the TXS 0218+357 QAS is produced by a laboratory amorphous olivine silicate template with a  measured peak optical depth of $\\tau_{10}=0.49\\pm0.02$. This relatively high optical depth is not a product of our adopted best-fitting template shape. All of the mineral fits over both the 10~$\\micron$ feature and combined 10 and 18~$\\micron$ features with $\\chi_r^2\\sim$1 result in a measured $\\tau_{10} \\gtrsim 0.45$. For the two mineral pairings, the combined 10~$\\micron$ peak optical depths are also consistently high. This derived peak optical depth for the TXS 0218+357 QAS is significantly larger than those measured for any of the 6 previously studied QASs, as illustrated in Table~\\ref{COMPARE}. Despite the deeper absorption, however, the best-fitting mineral species is consistent with several other QASs and with silicate dust grains in the Milky Way ISM. We discuss the implications of this high level of absorption and the dust grain species in the following sections.  In comparing the peak optical depth measurements from the different QASs we now use both the measured peak optical depth values ($\\tau_{10,obs}$), as well as the optical depths corrected ($\\tau_{10,cor}$) by accounting for the covering factor ($C_f$). This correction has been implemented using the expression \\begin{equation} \\tau_{cor}=-\\ln \\left( 1 - {1-e^{-\\tau_{obs}} \\over C_f} \\right) \\end{equation} The covering factor quantifies the fraction of the subtended area along the sightline to the continuum source obscured by the dusty material, and in many systems may be less than 1.0, thus implying a higher physical optical depth than that measured. Since direct estimates of the dust covering factor along the sightline are not feasible, we assume that the dust traces either the molecular or H~I gas, and has a similar covering factor. For the TXS 0218+357 QAS, both molecular line analyses [e.g., \\citet{Muller07}] and H~I studies [e.g., \\citet{Patnaik95}] find a difference between the measured fluxes for the southwest and northeast lensed quasar images, with a flux ratio of $4.2^{+1.8}_{-1.0}$ at 106 GHz \\citep{Muller07}. However, the precise value of the gas covering factor is debated in the literature [e.g., \\citet{Jethava07,Zeiger10}], and the value at 10~$\\micron$ may differ from that at $\\sim$100~GHz. It has also been suggested that the H~I absorption regions may be more extended than the region producing the molecular lines \\citep{Menten96}. For this analysis we adopt a value for the covering factor of $C_f=0.8$, which is broadly consistent with values found by \\citet{Wiklind95} for molecular material in this system. This results in a corrected optical depth value of $\\tau_{10,cor}=0.67\\pm0.04$ for the TXS 0218+357 QAS. For the z=0.886 QAS toward the lensed quasar PKS 1830-211, we adopt a lower total covering factor of $C_f=0.35$ \\citep{Frye97,Henkel08}, which raises the optical depth to $\\tau_{10,cor}=1.12\\pm0.04$. However, the value of $C_f$ is relatively uncertain [e.g., \\citet{Henkel08}], and a slight difference in the covering factor could have a significant impact. For instance, $C_f=0.3$ and $C_f=0.4$ would imply $\\tau_{10,cor}=1.54\\pm0.07$ and $\\tau_{10,cor}=0.89\\pm0.03$, respectively. For the non-lensed systems, covering factors may also be applicable. For AO 0235+164 we use $C_f=1.0$ \\citep{Curran05,Kanekar09}. For the Q0852+3435, Q0937+5628, and Q1203+0634 QASs estimates of the covering factor are not available in the literature. Although \\citet{Kanekar09} find a median covering factor of $C_f=0.885$ for a sample of z$<$0.9 DLAs, this value was determined at low frequencies ($\\lesssim$ 1.1~GHz) where extended radio structure may be contributing to the emission. Since such extended structure is unlikely to be associated with the silicate dust, we instead conservatively adopt a value of $C_f=1.0$ for all QASs without a published covering factor. For the z=0.437 QAS toward 3C~196, although the 21 cm absorption has a very low covering factor against the AGN [$C_f \\lesssim 0.1$; \\citet{Kanekar09}], it is not clear that such a low covering factor is appropriate for the dusty QAS material, due to the differences between the radio and optical sightlines to the quasar. Thus, we again adopt a representative $C_f=1.0$ for the 3C~196 QAS. Given the significant uncertainties and assumptions associated with some of these covering factors, when examining correlations with other dust and gas properties we utilize both the measured and the corrected peak optical depth measurements, rather than relying solely on the covering factor-corrected values.  \\begin{deluxetable}{llllllllll} \\tabletypesize{\\scriptsize} \\rotate \\tablecaption{Comparison with Other Program QASs} \\tablewidth{0pt} \\tablehead{ \\colhead{QAS} & \\colhead{z$_{abs}$} & \\colhead{Mineral} & \\colhead{$\\tau_{10,obs}$} &  \\colhead{$C_f$} & \\colhead{$\\tau_{10,cor}$} & \\colhead{E(B-V)} & \\colhead{EW$_{2796}$} & \\colhead{$\\tau_{21}$}& \\colhead{$\\tau_{21,cor}$}} \\startdata AO 0235+164 & 0.524 & Amorph.Oliv.GS (1)& 0.086$\\pm$0.003\\tablenotemark{a}~(2) & 1.0~(3,4) & 0.086$\\pm$0.003 & 0.227$\\pm$0.003 (5) & 2.42$\\pm$0.06 (6) & $\\geq$0.71 (7) &0.71\\\\ 3C 196 & 0.437 & Amorph.Oliv.2.0 (8) & 0.103$\\pm$0.007\\tablenotemark{a,b}~(2) & 1.0\\tablenotemark{c}& 0.103$\\pm$0.007 &0.20$\\pm$0.07 (2,9) & 2.00$\\pm$0.07 (10) & 0.05 (2,11)&\\nodata\\\\ Q0852+3435 & 1.309 & Amorph.Oliv.GPC (1)  &0.165$\\pm$0.010\\tablenotemark{a}~(2) &  1.0\\tablenotemark{c}& 0.165$\\pm$0.010 &0.40$\\pm$0.12 (2)& 2.97$\\pm$0.12 (2) &0.14 (2,12)&0.14\\\\ Q0937+5628 & 0.978 & Amorph.Oliv.GPC (1) & 0.192$^{+0.018}_{-0.017}$\\tablenotemark{a}~(2) & 1.0\\tablenotemark{c} &0.192$^{+0.018}_{-0.017}$ &0.34$\\pm$0.12 (2) & 5.77$\\pm$0.22 (2) & \\nodata&\\nodata\\\\ Q1203+0634 & 0.862 & Amorph.Pyrox (13) & 0.277$^{+0.041}_{-0.040}$\\tablenotemark{a,d}~(2) &1.0\\tablenotemark{c} & 0.277$^{+0.041}_{-0.040}$ &0.54$\\pm$0.18 (2) & 5.52$\\pm$0.30 (2)&\\nodata &\\nodata\\\\ PKS 1830-211 & 0.886 & Hortonolite (14) & 0.269$\\pm$0.006\\tablenotemark{e}~(15) & 0.35~(16,17) & 1.120$^{+0.041}_{-0.039}$ &0.57$\\pm$0.13(18)&\\nodata& 0.055 (19) &0.17\\\\ TXS 0218+357 & 0.685 & Amorph.Oliv.GPC (1) & 0.493$\\pm$0.024 & 0.8~(20,21) & 0.667$^{+0.036}_{-0.035}$ &0.62$\\pm$0.04~(18)&3.0 (22) &0.048 (23) &0.06\\\\ \\enddata \\tablenotetext{a}{Quasar continuum normalized spectrum from \\citet{Kulkarni11} re-fitted using methodology described in \\S\\ref{results}.} \\tablenotetext{b}{If a power-law quasar continuum normalization were adopted, a shallower optical depth (as low as $\\tau_{10}\\sim0.05$) could be obtained, depending on the adopted silicate template.} \\tablenotetext{c}{For QASs with no measurements of the covering factor in the literature, a value of 1.0 is assumed.} \\tablenotetext{d}{The fit using the Am.Oliv. template described in \\citet{Kulkarni11} has a consistent peak optical depth, but slightly poorer fit-quality.} \\tablenotetext{e}{Uncertainty recalculated to match methodology described in \\S\\ref{results}.} \\tablecomments{Comparison of the silicate absorption in the 7 studied QASs (column 1). We list the (column 2) QAS redshift; (column 3) laboratory mineral chosen from \\citet{Aller12} Table 10 list which produces the best fit to the 10~$\\micron$ feature (and reference); (column 4) measured peak 10~$\\micron$ silicate optical depth; (column 5) adopted covering factor (and reference); (column 6) peak 10~$\\micron$ silicate optical depth corrected using adopted covering factor; (column 7) E(B-V) reddening; (column 8) QAS rest-frame equivalent width for the Mg~II $\\lambda$2796 absorption (in \\AA); (column 9) the 21-cm optical depth; and (column 10) the 21-cm optical depth corrected using the adopted covering factor. The reddening, Mg~II equivalent width, and 21-cm optical depth are obtained from the literature, with the reference given in parenthesis. When multiple reddening values were presented in a single analysis, the value with the lowest $\\chi^2$ was adopted. References: (1) \\citet{Chiar06} based on data from \\citet{Henning99}; (2) \\citet{Kulkarni11}; (3) \\citet{Curran05}; (4) \\citet{Kanekar09}; (5) \\citet{Junkar04}; (6) \\citet{Wolfe77}; (7) \\citet{Roberts76}; (8) \\citet{Dorschner95}; (9) \\citet{Gharanfoli07}; (10) \\citet{Foltz88}; (11) \\citet{Briggs01};  (12) \\citet{Srianand08}; (13)  \\citet{Spoon} based on data from \\citet{Fabian01}; (14) \\citet{Jaeger} with tabulated data provided by C. J\\\"{a}ger; (15) \\citet{Aller12}; (16) \\citet{Frye97}; (17) \\citet{Henkel08};  (18) \\citet{Falco99};  (19) \\citet{Chengalur99}; (20) \\citet{Wiklind95}; (21) \\citet{Muller07}; (22) \\citet{Browne93}; \\& (23) \\citet{Carilli93}}\\label{COMPARE} \\end{deluxetable}  \\subsection{10-to-18~$\\micron$ Silicate Dust Optical Depth Ratio} Silicate dust grains in the ISM produce 2 prominent spectral features, which we see in the TXS 0218+357 QAS: a 10~$\\micron$ and an 18~$\\micron$ feature. The former is produced from an Si-O stretching mode in the dust molecules, while the latter originates in an O-Si-O bending mode. The amount of absorption produced by these two features, i.e. the ratio of their peak optical depths, varies with changes in dust grain properties such as the distribution of grain shapes and sizes. Analyses of the ISM in Local Group galaxies [e.g., \\citet{Weingartner01}], as well as those in higher redshift systems such as AGN tori [e.g., \\citet{Hao05}], have utilized the observed ratio of these features to constrain dust grain properties.  In the TXS 0218+357 QAS, we find the peak optical depth ratio calculated from the observed flux ratios to be $\\tau_{10}/\\tau_{18}= 1.31\\pm0.48$. The best-fitting (lowest $\\chi_r^2$) amorphous olivine template, which had porous, ellipsoidal particles \\citep{Chiar06}, has a $\\tau_{10}/\\tau_{18}$ ratio of 1.41, which more closely matches the data than, for example, a similar amorphous olivine template with solid, spherical particles that predicts a $\\tau_{10}/\\tau_{18}$ ratio of 1.62. (However, it is not possible to definitively rule out solid, spherical particles given the large uncertainty in the  $\\tau_{10}/\\tau_{18}$ ratio inferred from the data.) While in principle, we could extend our analysis to compare our observed silicate dust opacity curves with models constructed for different grain temperatures and geometries [e.g. \\citet{Li01,Weingartner01}], we have not taken this step because of the low S/N of our data, and the uncertainties associated with our quasar continuum normalization. As discussed in Appendix~\\ref{normAP}, since our data do not extend past 21.5~$\\micron$ in the QAS rest-frame, the shape of the quasar continuum underlying the 18~$\\micron$ silicate feature is poorly constrained. Differing assumptions about the physical shape of the quasar continuum can significantly affect the depth of the 18~$\\micron$ feature. We, therefore, feel that a rigorous comparison with the dust grain models is not warranted for this dataset.  \\subsection{Connection between Reddening and Silicate Absorption}  Given the high peak optical depth measurement for the silicate dust in the TXS 0218+357 QAS relative to the 6 previously studied QASs, we have revisited the suggested trend between silicate absorption peak optical depth and reddening presented in \\citet{Kulkarni11}. In that work, it was noted that the trend seemed similar to that for Milky Way diffuse ISM clouds, i.e. $\\tau_{10}=0.17 \\times E(B-V)$ for R$_V$=3.1, but that the slope differed such that the ratio of $\\tau_{10}/E(B-V)$ was 2-4 times that of diffuse Galactic clouds. This was used to argue for either a greater abundance of silicate dust in these QASs than is found in the Milky Way ISM dust, possibly stemming from differences in the underlying stellar populations, or for differences in the dust grain geometrical and size properties which would affect R$_V$.  The addition of the TXS 0218+357 QAS and the z=0.886 PKS 1830-211 QAS to the 5 previously studied QASs reveals a clear trend of increasing reddening with increasing silicate measured peak optical depth, as illustrated in Figure~\\ref{Ebvplot} (left). The total reddening measurements are from the literature, as detailed in Table~\\ref{COMPARE}. Considering all 7 QAS we find a Pearson linear correlation coefficient of 0.89 (0.8\\% probability of no correlation), a Spearman (nonparametric) rank correlation coefficient of 0.89 (significance 0.007, i.e., 0.7\\% probability of no correlation), and a (nonparametric) Kendall's Tau value of 0.71 (significance 0.02). The small probabilities/significances indicate a relatively strong correlation. We fitted a linear least-squares relationship to the 7 QASs, weighting by both the uncertainties in the reddening from the literature and by those for our measured peak silicate optical depths. We find a slope of 1.00$\\pm$0.11, which is substantially steeper than the 0.17 slope found for the extrapolation for Milky Way diffuse clouds from \\citet{Roche84}. Our fit is illustrated in Figure~\\ref{Ebvplot} (left) with a solid black line, while the shallower Milky Way extrapolation is shown as a blue dot-dashed line. All of the studied QASs lie significantly above the Milky Way extrapolation. We note that the steepness of the slope we obtain is being driven in part by the TXS 0218+357 QAS. However, even excluding that QAS gives a slope of 0.60$\\pm$0.18 (dashed red line), which is much steeper than the Milky Way extrapolation. If we require our relationship to extrapolate as a smooth linear function to dust-free systems, i.e. to predict an optical depth of 0.0 for no reddening, we would get a still relatively steep slope of 0.43$\\pm$0.01 (shown as dotted orange line in Figure~\\ref{Ebvplot}). While this slope is closer to that predicted from the Milky Way, it remains nearly three times higher than the slope for the extrapolated Milky Way relation, and the TXS 0218+357 QAS would lie significantly above the relationship. Since the reddening error bars are sensitive to the adopted intrinsic quasar spectrum in the extinction-curve-fitting reddening measurements [e.g., \\citet{Kulkarni11}] we have also examined the slope obtained if we omit the reddening uncertainties when fitting over all 7 QASs. We find a slope of 0.51$\\pm$0.02, which does not alter our conclusions.  We have also explored the relationship between the reddening and the silicate peak optical depth using the covering factor-corrected peak optical depths, as depicted in Figure~\\ref{Ebvplot} (right). Since the covering factors for several of the QASs are highly uncertain, we place less confidence in the derived slope of the relationship, although there is still a general trend of increasing silicate peak optical depth with increasing reddening. Furthermore, since the optical depth values are increased by applying the covering factor correction, the offset relative to the Milky Way diffuse cloud extrapolation increases for $C_f<1$. The Pearson correlation coefficient is 0.76 (5\\% probability of no correlation), with a Spearman (nonparametric) rank correlation coefficient of 0.89 (significance 0.007), and a (nonparametric) Kendall's Tau value of 0.71 (significance 0.02). The magenta dash-dot-dot-dotted line depicts the linear least-squares fit (slope of 0.68$\\pm$0.31) obtained when excluding the TXS 0218+357 and PKS 1830-211 QASs. When corrected for the covering factors, the molecule-rich PKS 1830-211 and TXS 0218+357 QASs lie above the trend suggested by the other QASs, i.e., their silicate peak optical depths are higher than expected given their reddening values. This is particularly interesting in the context of studies in the Milky Way [e.g., \\citet{Chiar07}], which find that while there is a tight correlation between the 10~$\\micron$ silicate peak optical depth and the near-infrared color excess $E(J-K)$ along diffuse ISM sightlines, when molecular clouds are examined the correlation is absent, with molecular sightlines exhibiting \\textit{lower} silicate optical depths for a given $E(J-K)$. The reason for this difference between diffuse and molecular cloud sightlines is still not completely clear. It has been argued that the origin is an increase in $E(J-K)$ in molecular clouds, possibly linked to grain growth or variations in dust species abundances \\citep{vanBreemen11}. Our molecule-rich systems, however, lie above the trend established by the other QASs. This is more similar to the Galactic center sightline, which exhibits a higher silicate optical depth for its observed $E(J-K)$ relative to other Galactic diffuse ISM sightlines, possibly because of fewer C-rich stars near the Galactic center. \\citep{Roche85}.  Overall, the disparity between the Milky Way extrapolation and the studied QASs seems most significant in systems with large silicate optical depths; these systems drive the steeper derived slope relative to the Milky Way. This may indicate that in distant galaxies at relatively high silicate dust columns, the selective extinction differs from that in the Milky Way because of differing dust grain properties. One possible explanation is that the interstellar dust grains may be larger than those in the Milky Way ISM. Large dust grains would yield a low UV extinction, which would translate into a high R$_V$, consistent with the measured R$_V\\sim$7.2 for the TXS 0218+357 QAS \\citep{Falco99}. This decrease in UV extinction, which is not accompanied by a similar decrease in IR absorption, could produce the offset relative to the Milky Way extrapolation. Furthermore, we note that reddening measurements in the Milky Way sample an integrated combination of different dust grain populations along the Galactic plane sightline, while systems such as the TXS 0218+357 QAS are viewed nearly face-on \\citep{York05}. This difference in viewing angle may mean that a single, less common, type of dust grain could disproportionately impact the derived total reddening for face-on galaxies. However, the good agreement between our silicate absorption profile and that for amorphous olivine, which has also been used to describe Milky Way ISM silicate dust grains, would imply that the silicate dust grains are somewhat similar. This might suggest that the disparity stems instead from a difference in the underlying stellar populations producing the dust. A larger sample of QASs, particularly QASs with large silicate optical depths, are needed to more strongly constrain the slope of the reddening-silicate optical depth relationship, and to investigate whether the relationship is adequately described by a simple linear trend. Furthermore, measurements of the silicate optical depth in QASs with low reddening values, in which the silicate optical depth is expected to be low requiring (future) high-sensitivity instruments, will help to establish whether the trend extends to systems with less dust.  \\begin{figure} \\epsscale{1.0} \\plottwo{Figure5a.eps}{Figure5b.eps}% \\caption{Illustration of possible correlation between the silicate peak optical depth and the reddening. The left panel utilizes the measured 10~$\\micron$ peak optical depths, while in the right panel the optical depths are corrected using the covering factor. In the left panel, the solid black line shows the linear least-squares fit to all 7 QASs, and the red dashed line shows the fit which excludes the TXS 0218+357 absorber. In the right panel, the magenta dash-dot-dot-dotted line shows the fit excluding both the TXS 0218+357 and PKS 1830-211 QASs. In both panels, the orange dotted line shows the fit required to pass through the origin. The shallow blue dot-dashed line is an extrapolation of the trend for Milky Way diffuse clouds from \\citet{Roche84}. [See the electronic edition of the Journal for a color version of this figure.]} \\label{Ebvplot} \\end{figure}  \\subsection{21-cm Absorption in TXS 0218+357} The 21-cm peak optical depth provides an estimate of the absorption produced by the cold, dusty H~I gas in the ISM. We list measurements of the observed $\\tau_{21cm}$, taken from the literature, in Table~\\ref{COMPARE}. In the TXS 0218+357 system, \\citet{Carilli93} find a peak optical depth of $\\tau_{21cm}=0.048\\pm0.005$, assuming a covering factor of 1.0. The line structure is noted to be somewhat asymmetrical, but under an assumed Gaussian fit a FWHM of 43$\\pm$4 km/s is derived. The velocity width, optical depth, and shape for this 21-cm absorption feature are all noted by  \\citet{Carilli93} to be similar to those seen in the z=0.437 QAS toward 3C~196, although concerns that the 3C~196 absorption is occurring against the lobes rather than against the AGN core imply that they may not be directly comparable \\citep{Briggs01}. This sightline difference may explain why despite the similarity in their measured 21-cm peak optical depths, the measured peak silicate optical depths differ: $\\tau_{10}$ is nearly five times larger for the TXS 0218+357 QAS than for the z=0.437 QAS toward 3C~196. The remaining systems with 21cm absorption measurements suggest that there may be a slight anti-correlation between the measured silicate peak optical depth and the measured 21 cm peak optical depth, as depicted in Figure~\\ref{T21plot} (left). However, as evidenced by the poor match of the linear least-squares fit to the data, this may not be a purely linear relationship, but more data points are required to establish the precise shape of the relationship. When the covering factors are applied to the 21 cm and silicate peak optical depth measurements (Figure~\\ref{T21plot}, left), the trend becomes less clear, largely because of the high derived $\\tau_{10,cor}$ optical depth for the PKS 1830-211 system, which has a low covering factor. Although a slight anti-correlation is suggested by the data, a larger data sample is required to investigate further. In any case, variations in properties of the material, such as the dust-to-gas ratio and the depletions, could induce scatter in even an intrinsically tight correlation.  \\begin{figure} \\epsscale{1.0} \\plottwo{Figure6a.eps}{Figure6b.eps}% \\caption{21 cm H~I optical depth plotted versus the silicate peak optical depth for the 5 QASs with published 21 cm measurements. The left panel utilizes the measured peak optical depths, while in the right panel the optical depths are corrected using the covering factor. The 3C 196 QAS is depicted in green in the left panel, and excluded in the right panel, since the optical and radio sightlines may differ in this system, and a reliable estimate of the covering factor is not available for this QAS. In the left panel, the black solid line shows a linear fit using all 5 QASs, while the red dashed line excludes the 3C 196 QAS measurement from the fit. There appears to be a suggestion of a slight anti-correlation between the 10~$\\micron$ and 21 cm peak optical depths. [See the electronic edition of the Journal for a color version of this figure.]} \\label{T21plot} \\end{figure}  \\subsection{Correlation between Mg~II Equivalent Width and Silicate Dust} We also investigate suggestions of a correlation between the Mg~II $\\lambda$2796 rest-frame equivalent width and the peak optical depth of the silicate dust absorption. \\citet{Kulkarni11} found a correlation between these properties and suggested that the more silicate-rich QASs may be more massive. This is because since the Mg~II lines are saturated in the absorbers studied by \\citet{Kulkarni11}, the Mg~II equivalent width is a better proxy for the velocity spread (i.e. the absorber mass or presence of outflows into the CGM) than for the elemental abundance of metals in the ISM gas.  We have revisited the  correlation here with the addition of the TXS 0218+357 QAS, as shown in Figure~\\ref{Mgplot} using both the measured (left) and corrected (right) silicate peak optical depths. However, we find that the TXS 0218+357 QAS is an outlier relative to the other 5 QASs, with a much stronger peak optical depth than would be predicted from its Mg~II equivalent width. Including all 6 QASs and using the measured peak optical depths, these parameters have a Pearson correlation coefficient of 0.22 (68\\% probability of no correlation), a Spearman correlation coefficient of 0.71 (significance 0.11), and a Kendall's Tau value of 0.47 (significance 0.19). If we exclude the TXS 0218+357 outlier point, however, we would obtain a stronger Pearson correlation coefficient of 0.86 (6\\% probability of no correlation) and Kendall's Tau of 0.60 (significance 0.14). A linear least-squares fit to all 6 QASs, weighting solely by the $\\tau_{10}$ uncertainties, since measurement uncertainties are not available for all of the Mg~II measurements in the literature, produces a slope of 0.042$\\pm$0.005 for the measured optical depth values, and 0.040$\\pm$0.005 using the covering-factor corrected optical depths. This relationship is illustrated as a dashed line. The previously studied 5 QASs cluster around this line, while the TXS 0218+357 QAS lies substantially above the fit.  If the TXS 0218+357 QAS is excluded from the fit, the slope decreases slightly to 0.035$\\pm$0.005 (same for the corrected optical depths), as depicted with the dotted line. A larger sample will be required to see whether the TXS 0218+357 QAS is anomalous, or whether the suggested correlation is not as tight was was implied with the earlier sample of QASs.  Some of the discrepancy for the TXS 0218+357 QAS may arise if the Mg~II line in this system is not saturated, unlike in the 5 previously studied QASs. This may be caused in part by the Mg~II slit encompassing both quasar sightlines, one of which has significantly less absorption and is brighter at optical wavelengths. While the equivalent width ratio of the lines is suggestive of saturation, the resolution and S/N of the spectrum are not high enough to definitively assess if the lines are saturated \\citep{Browne93}. If the lines are unsaturated, the Mg~II measurement may be providing an estimate more reflective of the elemental abundance than of the galaxy dynamics (i.e. the velocity spread). Alternatively, it is possible that the sampled sightline through the TXS 0218+357 QAS's nearly face-on galaxy probes inner regions which produce less efficient galactic outflows of warm gas containing Mg~II, and that if viewed from a different sightline a higher Mg~II equivalent width would be obtained. If an elemental discrepancy is invoked, the ISM gas in this system may be Mg-poor either because of an intrinsic deficit in the total (gas+solid phase) amount of Mg, which would also result in less Mg-rich (more Fe-rich) species of olivines/pyroxene silicates. Alternatively, the gas may be Mg-poor because of depletions onto the dust grains, resulting in Mg-rich olivines/pyroxenes. Data with higher spectral resolution and S/N, and with longer wavelength coverage, are required to discriminate better between these silicate species.  \\begin{figure} \\epsscale{1.0} \\plottwo{Figure7a.eps}{Figure7b.eps}% \\caption{Relation between the rest-fame Mg~II $\\lambda$2796 equivalent width and silicate peak optical depth for the 6 QASs with published Mg~II measurements. The left panel utilizes the measured 10~$\\micron$ peak optical depths, while in the right panel the optical depths are corrected using the covering factor. The black dashed line shows the fit using all 6 QASs, while the red dotted line depicts the fit excluding the TXS 0218+357 QAS. The TXS 0218+357 QAS appears to be an outlier compared to the relationship between these parameters for the other program QASs. [See the electronic edition of the Journal for a color version of this figure.]} \\label{Mgplot} \\end{figure}  \\subsection{Connections Between Silicate Crystallinity and Molecule Formation?} Lastly, the TXS 0218+357 QAS provides an interesting opportunity to investigate whether there is a link between silicate crystallinity and a high enrichment of molecular material. The degree of silicate crystallinity depends on the conditions (e.g., temperature) under which the dust grains formed and the subsequent processing of the dust grains (e.g., by cosmic rays, grain-grain collisions), and, therefore, has implications for the evolution of the dust grains [e.g., \\citet{Kemper11}]. In the z=0.886 QAS toward PKS 1830-211, a system which has been noted to have the largest number (28) of detected molecular species in any extragalactic object to date \\citep{Muller11}, we previously found very suggestive evidence for a high degree of silicate crystallinity \\citep{Aller12}. This evidence for crystallinity stems from the shape and location of the substructure in the observed QAS 10~$\\micron$ silicate absorption feature. Since in the Milky Way ISM the silicate dust is predominately amorphous \\citep{Kemper04, Li07}, this raises the question of how the silicate dust in PKS 1830-211 could be so much more crystalline, and whether it is linked to the high degree of molecule-enrichment. For instance, does some aspect (e.g. temperature, density) of the ISM environment facilitate the formation of both molecules and crystalline silicate grains? Since the TXS 0218+357 QAS is also noted to be very rich in molecules, and exhibits 10~$\\micron$ silicate absorption, it provides a unique opportunity to test this hypothesis.  The TXS 0218+357 QAS has been noted to exhibit signatures of many molecules, including CO, HCO$^+$, HCN, HNC, CN, CS, C$^{18}$O, $^{13}$CO, H$_2$O, NH$_3$, OH, H$_2$CO, and LiH \\citep{Wiklind95, Combes95, Combes97b, Combes98,Combes97,Gerin97,Menten96,Kanekar02,Kanekar03,Henkel05,Jethava07,Zeiger10,Friedel11}. As for the lensed PKS 1830-211 z=0.886 QAS, it appears that one of the two quasar sightlines (both of which are encompassed by the \\textit{IRS} slit) is more enriched. Although a direct comparison of the abundances in the z=0.886 PKS 1830-211 QAS and the TXS 0218+357 QAS is challenging, since derived molecular abundances depend on assumptions about the degree of line saturation, the covering factor, and the excitation temperature, for most molecules it appears that the TXS 0218+357 QAS is slightly less enriched. The OH and HCO$^+$ column densities are both good predictors of N$_{H_2}$ in Galactic clouds \\citep{Liszt99}. A direct comparison in \\citet{Kanekar02} finds that the derived N$_{OH}$ and N$_{HCO+}$ column densities in TXS 0218+357 are 19-23\\% of those for the z=0.886 PKS 1830-211 absorber. Using OH as a tracer of the molecular hydrogen results in a column density of N$_{H_2}$ of $2.65\\times 10^{22}$ cm$^{-2}$ for the TXS 0218+357 \\citep{Kanekar02}.  However, while the evidence for significant molecule enrichment in the ISM in the TXS 0218+357 QAS is undisputed, the evidence for or against silicate crystallinity is much less clear due to the low S/N of our data. The combination of a lower flux density, fewer exposures, and a smaller total exposure time for the TXS 0218+357 quasar, in comparison with PKS 1830-211 \\citep{Aller12}, makes it much more difficult to establish the presence or absence of silicate crystallinity. As discussed in \\S\\ref{results} the laboratory amorphous olivine laboratory template for the TXS 0218+357 system provides a good fit, both in terms of the breadth and depth of the 10 and 18~$\\micron$ silicate absorption features, and in terms of the peak absorption wavelength. However, as illustrated in Figure~\\ref{SINGLEFIT}, the magnitude of the substructure in the 10 and 18~$\\micron$ regions for some species of crystalline olivine and pyroxene are small enough to be masked by the large uncertainties and noise associated with our spectrum. Several of the predicted substucture features align relatively well with the possibly noise features present in our data. Furthermore, as established with the two-mineral fits in \\S\\ref{twofitsect}, a combination of crystalline olivine and pyroxene would produce an excellent match to the 10~$\\micron$ and 18~$\\micron$ silicate absorption features; the predicted abundant substructure would occur at levels below that of the noise in the data. Thus, although the smooth, featureless amorphous olivine fit is both consistent with the ISM seen in the Milky Way and with our data, we cannot rule out that abundant silicate crystallinity \\textit{could} be present in the TXS 0218+357, although it seems less likely than in the PKS 1830-211 z=0.886 QAS.  If the explanation that the TXS 0218+357 QAS is produced by amorphous olivines is adopted, it establishes a similarity between the silicate dust in this z=0.685 galaxy and that  in the Milky Way ISM. By contrast, independent of the substructure, the redder peak wavelength of the absorption feature in the PKS 1830-211 z=0.886 QAS implies silicate dust grain differences with that in the Milky Way ISM and with that in the TXS 0218+357 QAS ISM. In combination, these data suggest that there could be significant variations in the dust grain properties between different galaxies, which may not be a simple function of redshift evolution or of molecular gas abundance. Further investigation of the silicate dust grains in the TXS 0218+357 with higher spectral resolution and S/N data, as well as a study of a broader sample of QASs, is required to quantify these variations and establish whether ISM silicate dust grain crystallinity varies between galaxies as a function of composition, star-formation rate, or redshift. Such studies may be possible within the next decade using new facilities such as \\textit{SOFIA} and \\textit{JWST-MIRI}."
    },
    "1402/1402.4812_arXiv.txt": {
        "abstract": "We consider dust formation during the ejection of the hydrogen envelope of a red supergiant during a failed supernova (SN) creating a black hole.  While the dense, slow moving ejecta are very efficient at forming dust, only the very last phases of the predicted visual transient will be obscured. The net grain production consists of $M_d \\sim 10^{-2} M_\\odot$ of very large grains ($10$ to $1000\\mu$m).  This means that failed SNe could be the source of the very large extrasolar dust grains identified by {\\it Ulysses}, {\\it Galileo} and radar studies of meteoroid re-entry trails rather than their coming from an ejection process associated with protoplanetary or other disks. ",
        "introduction": "\\label{sec:introduction}  The formation of black holes during a failed supernova (SN) has generally been assumed to lead to a black hole with the mass of the star at the time core collapse is initiated (e.g. \\citealt{Heger2003}). \\cite{Nadezhin1980} argued that this was not true of red supergiants because the hydrogen envelope is so weakly bound that the drop in the gravitational potential created by neutrino losses is enough to unbind the envelope.  Radiation hydrodynamic simulations by \\cite{Lovegrove2013} confirmed this supposition.  Failed SNe of red supergiants lead to the formation of black holes with (roughly) the mass of the helium core of the star at the time of collapse and not the total mass of the star.  This new expectation is of particular interest because studies of SN progenitors appear to be finding too few massive progenitors (\\citealt{Kochanek2008}) particularly in the upper end of the red supergiant mass range (\\citealt{Smartt2009}). \\cite{Smartt2009} estimated that red supergiant progenitors were missing from $(16.5\\pm1.5)M_\\odot$ to the upper mass limit for stars to explode as red supergiants and become Type~IIP SNe ($25$-$30M_\\odot$).  This is an interesting mass range for failed SNe because theoretical studies indicate that many stellar models in this mass range have density structures that make it more difficult to produce a successful explosion (e.g. \\citealt{Oconnor2011}, \\citealt{Ugliano2012}). Other options include changing stellar models so that stars in this mass range do not end their lives as red supergiants (e.g. \\citealt{Groh2013}) or using dusty winds to make them less observable (e.g. \\citealt{Walmswell2012}, but see \\citealt{Kochanek2012}).  Producing failed SNe in this mass range also provides a natural explanation for the observed distribution of remnant masses (\\citealt{Kochanek2013}).  The remnant mass function is bimodal between neutron stars with masses near $1.4M_\\odot$ and black holes with masses of $5$-$10M_\\odot$, although it is uncertain if the gap in the mass function is truly empty or simply a deep minimum (e.g. \\citealt{Bailyn1998}, \\citealt{Ozel2010}, \\citealt{Farr2011}, \\citealt{Kreidberg2012}). This gap is difficult to explain if black holes form by fall back of material in a successful explosion (e.g. \\citealt{Zhang2008}, \\citealt{Fryer2012}). The gap is, however, a natural consequence of combining supernovae without fall back, which best explain the masses observed in neutron star binaries (\\citealt{Pejcha2012}), with a population of $\\sim 20M_\\odot$ failed supernovae where the typical black hole mass of $\\sim 7M_\\odot$ is simply the typical helium core mass of the failed SNe.  \\cite{Lovegrove2013} also pointed out that the ejection of the envelope would be associated with a relatively long lived ($\\sim$ year), modest luminosity ($\\sim 10^6 L_\\odot$) transient largely powered by the recombination of the envelope.  We have been carrying out a search for failed SNe using the Large Binocular Telescope following the ideas of \\cite{Kochanek2008}, where we monitor nearby galaxies to see if any stars effectively vanish, potentially with some intervening transient phenomenon.  This survey can easily identify the transients predicted by \\cite{Lovegrove2013}, although it is not well suited for identifying the shorter ($\\sim$~week) shock break out peak associated with the event (see \\citealt{Piro2013}).  These more massive red supergiants tend to have winds that produce silicate dusts and other products of oxygen-rich chemistry (see, e.g., the review by \\citealt{Cherchneff2013}), and, as mentioned earlier, \\cite{Walmswell2012} even propose using this obscuration as a possible means of solving the red supergiant problem identified by \\cite{Smartt2009}.  During the envelope ejection following a failed SNe, this same material is going to be ejected {\\it en masse} at much higher densities and only moderately higher velocities, so presumably it is also likely to form dust but in much greater quantities.  This leads to the possibility that the optical properties of the transient are greatly modified from the dust free simulations carried out by \\cite{Lovegrove2013}.  We already know of one class of transients that successfully cloaks themselves in dust following their peaks, the SN~2008S class of explosive transients from AGB stars (see, e.g., \\citealt{Thompson2009}, \\citealt{Kochanek2011}, \\citealt{Szczygiel2012}). The mechanism here is very different, since the dust in the SN~2008S class is almost certainly dust re-forming in a pre-existing dense wind after being destroyed by the transient (\\citealt{Kochanek2011}).  However, it does suggest an investigation of whether the visual transient predicted by \\cite{Lovegrove2013} will be substantially modified by dust formation and if the ejected dust has any consequences for our understanding of the interstellar medium. In \\S2 we model dust formation in these systems to find that the transients are not substantially modified even though they should be very efficient at forming significant masses of very large dust grains.  We discuss the implications in \\S3. ",
        "conclusions": "Despite having very favorable conditions for dust formation, the optical signature of a failed SNe from a red supergiant will be little affected by dust formation.  Failed SNe are, however, very efficient at producing dust compared to successful SNe.  Keeping all else fixed, the final grain size in Equation~\\ref{eqn:afinal} is $a_f \\propto (v L_*)^{-1}$ because the growth rate at the dust formation radius scales as $da/dt \\propto R_f^{-3} \\propto L_*^{-3/2}$ while the time scale over which the growth occurs is $R_f/v \\simeq L_*^{1/2}/v$. Since for a failed SNe $v \\simeq 10^2$~km/s and $L_* \\simeq 10^6 L_\\odot$ and for a successful SNe $v \\simeq 5000$~km/s and $L_*\\simeq 10^9 L_\\odot$, the nominal final size ratio of $a_{fail}/a_{succ} \\sim 50000$ gives a sense of the difference.  The hotter radiation temperature of successful SNe further favors the failed SNe.  The one advantage for the successful SNe is that the heavily metal enriched material outside the iron core is also ejected, so some fraction of the ejecta from a SNe has a condensible fraction of $X \\sim 10^{-1}$ instead of $X \\sim 10^{-3}$ for the envelope of a red supergiant. Somewhat surprisingly, even the very metal rich ejecta of real SNe have slower dust growth rates than the envelope of a failed SN.  The conditions are so favorable for dust formation that the grains formed in this scenario are enormous -- up to $a_f \\sim 1$~mm in size. This is interesting because several recent experiments have found evidence for a significant population of very large extrasolar dust grains.  Impact detectors on the {\\it Ulysses} and {\\it Galileo} spacecraft (\\citealt{Landgraf2000}, \\citealt{Kruger2007}) measured a significant flux of large grains entering the heliosphere up to their detection limit of $a \\simeq 1\\mu$m, and still larger 10-30$\\mu$m grains of apparently extrasolar origin have been found by ground-based radar studies of meteoroid re-entry trails (e.g. \\citealt{Meisel2002}, \\citealt{Baggaley2000}). The resulting rates appear to be consistent with simply extending the roughly logarithmic distribution of grain masses, $m(dn/dm) \\propto m^{-1.1}$, (e.g. \\citealt{Murray2004}, \\citealt{Draine2009}) well beyond the standard truncation scale of the models for dust in the interstellar medium (ISM) at $a \\sim 0.1\\mu$m.  There has been no natural source of these very large grains because even the most extreme stellar winds have natural size cutoffs of order $1\\mu$m.  There there are grains with AGB star compositions as large as $10\\mu$m found in meteorites, but they are rare (\\citealt{Zinner1998}).  For still larger grains, ejection from proto-planetary disks may be possible (see \\citealt{Murray2004}). These failed SNe would produce $\\sim 10\\%$ of all dust mass (assuming SN and stars produce comparable amounts of dust) with a mechanism that both naturally creates and ejects the grains into the interstellar medium (ISM) at relatively low velocities so as to minimize sputtering.  Quantitatively, this would be too little mass to explain the solar system results.  However, there is almost certainly a problem with the solar system abundance normalizations because simply extending the ISM size distribution (e.g. \\citealt{Draine2003}) to larger sizes would lead to an abnormal average extinction curve ($R_V \\simeq 5.8$) and require a greater abundance of condensible species (C, O, Fe, etc.) than is available (\\citealt{Draine2009}). Possible solutions include a local anomaly in the grain size distribution (\\citealt{Draine2009}), propagation effects from the ISM into the solar system (\\citealt{Slavin2012}) or that the source of the grains is not actually extrasolar (\\citealt{Belyaev2010}). The small dust mass fraction from failed SNe would not create these problems.  \\cite{Socrates2009} discuss a means of directly detecting large grains in the ISM as a scattered light halo during optical transients.  At present, however, it seems that interest in these large grains has waned -- a relation to black hole formation, however bizarre, might provide a motivation to resolve these issues."
    },
    "1402/1402.2480_arXiv.txt": {
        "abstract": "{Asteroseismology may in principle be used to detect unresolved stellar binary systems comprised of solar-type stars and/or red giants. This novel method relies on the detection of the presence of two solar-like oscillation spectra in the frequency spectrum of a single lightcurve. Here, we make predictions of the numbers of systems that may be detectable in data already collected by the NASA \\emph{Kepler} Mission.} Our predictions, which are based upon TRILEGAL and BiSEPS simulations of the \\emph{Kepler} field of view, indicate that as many as 200 or more ``asteroseismic binaries'' may be detectable in this manner. Most of these binaries should be comprised of two He-core-burning red giants. Owing largely to the limited numbers of targets with the requisite short-cadence \\emph{Kepler} data, we expect only a small number of detected binaries containing solar-type stars.  The predicted yield of detections is sensitive to the assumed initial mass ratio distribution of the binary components and therefore represents a sensitive calibration of the much debated initial mass ratio distribution near mass ratio unity. ",
        "introduction": "\\label{sec:intro}  The NASA \\emph{Kepler} Mission has provided a wealth of data for studying stellar binary systems. More than 2000 eclipsing binaries have so far been found \\citep{Prsa11, Slawson11}. Circumbinary planets have also been discovered in several of these systems (e.g., see \\citealt{Welsh12, Orosz12}). The \\emph{Kepler} eclipsing binary catalogue\\footnote{see {\\tt http://keplerebs.villanova.edu/}} provides information on binaries with periods up to a few hundred days.  In this paper we discuss the prospects for using asteroseismology as a novel way of discovering non-eclipsing binary systems in numbers that would potentially provide a useful statistical sample, out to periods much longer than the eclipsing binary catalogue.  Asteroseismology has been one of the major successes of \\emph{Kepler}. It has in particular revolutionized the study of stars that show solar-like oscillations, with detections made in several hundred cool main-sequence and subgiant stars \\citep{Chaplin11a, Chaplin14}, and over ten-thousand red giants (e.g., see \\citealt{Bedding11, Hekker11, Huber11, Huber13a}). These stars present rich spectra of overtones in the frequency-power spectrum.  Asteroseismology may be used to discover candidate binaries by detecting the presence of two solar-like oscillation spectra in the frequency spectrum of a single \\emph{Kepler} lightcurve. The oscillation spectra may be well separated in frequency; or they may overlap, with the individual oscillation peaks of the two stars interspersed. {To illustrate the concept, we have created an artificial example by using \\emph{Kepler} data on two red giants in the open cluster NGC\\,6819. We combined the lightcurves of two cluster members -- KIC\\,5024405 and KIC\\,5112950 -- with detected oscillations (from Fig.~2 of~\\citealt{Stello10}).  The bottom of our Fig.~\\ref{fig:combined} shows the frequency-power spectrum of the combined lightcurve (i.e., the simulated asteroseismic binary), whilst the upper two panels show the spectra of the two original lightcurves. Arrows mark the locations of the solar-like oscillations.}  The analysis required to discriminate the constituent oscillation frequencies of stars comprising asteroseismic binaries is beyond the intended scope of the present paper, although we add that in some scenarios -- when the oscillations spectra are well separated in frequency -- it is a fairly trivial matter. {That would have been the case for our artificial example: the oscillations are clearly separated (because KIC\\,5112950 is about 2.5-times more luminous than KIC\\,5024405).}  Our aim in this paper is to estimate -- subject to different assumptions about the underlying binary population -- the number of binaries in the \\emph{Kepler} field where both components would be expected to show detectable solar-like oscillations in \\emph{Kepler} data. Our results are based on predictions from two stellar population simulations of the \\emph{Kepler} field, to which we applied target selection criteria based on those used to select the real \\emph{Kepler} target list. We consider binaries comprised of solar-type stars and/or red giants. Imposing the seismic detectability constraint for both components strongly favours binaries with mass ratio near unity. The predicted number therefore represents a very sensitive calibration of the much debated initial mass ratio distribution (IMRD; e.g. \\citet{goodwin13} and references therein) near mass ratio unity.  The layout of our paper is as follows. In Section~\\ref{sec:popsyn} we discuss the population synthesis predictions made by TRILEGAL \\citep{Girardi05, Girardi12} (Section~\\ref{sec:tri}) and BiSEPS \\citep{Farmer13} (Section~\\ref{sec:biseps}); and target selection applied to the raw predictions (Section~\\ref{sec:targ}). In Section~\\ref{sec:det} we explain how we used the data on the intrinsic properties of the synthetic target-selected populations to make the asteroseismic detection predictions. Section~\\ref{sec:res} presents results on the asteroseismically detected synthetic binaries. We end in Section~\\ref{sec:disc} with concluding remarks on the implications of our results for analysis and exploitation of the real, available \\emph{Kepler} data. ",
        "conclusions": "} \\label{sec:disc}  We have used results from two stellar population simulations of the \\emph{Kepler} field of view to consider {potential detection yields} from applying a novel asteroseismic method to discover binary systems comprised of solar-type stars and/or red giants. The method relies on the detection of the presence of two solar-like oscillation spectra in the frequency spectrum of a single \\emph{Kepler} lightcurve. Our predictions suggest that 200 or more binaries may be detectable in this manner using the \\emph{Kepler} data. Most of these binaries should be comprised of two He-core-burning red giants. Owing largely to the limited numbers of targets with the requisite short-cadence \\emph{Kepler} data, we expect only a small number of detected binaries containing solar-type stars.  The method is not biased by, or dependent on, the inclination, period (separation) or velocities of the constituent components. As such, it may provide the only way to detect some of the binaries that are in the \\emph{Kepler} database. The predicted yield of detections is sensitive to the IMRD, which is a robust indicator of the pristine mass distribution at the epoch of star formation (e.g., see \\citealt{parker13}). Given the additional constraints, the detection of asteroseismic binaries will also provide targets that are useful for testing stellar models and will add significantly to the list of eclipsing binaries with a solar-like oscillating component \\citep[e.g., see][]{Hekker10, Gaulme13}.  {Work is now underway to discover asteroseismic binaries in the \\emph{Kepler} data archive.}"
    },
    "1402/1402.3679_arXiv.txt": {
        "abstract": "The surface winds of Jupiter and Saturn are primarily zonal. Each planet exhibits strong prograde equatorial flow flanked by multiple alternating zonal winds at higher latitudes. The depth to which these flows penetrate has long been debated and is still an unsolved problem. Previous rotating convection models that obtained multiple high latitude zonal jets comparable to those on the giant planets assumed an incompressible (Boussinesq) fluid, which is unrealistic for gas giant planets. Later models of compressible rotating convection obtained only few high latitude jets which were not amenable to scaling analysis.  Here we present 3-D numerical simulations of compressible convection in rapidly-rotating spherical shells. To explore the formation and scaling of high-latitude zonal jets, we consider models with a strong radial density variation and a range of Ekman numbers, while maintaining a zonal flow Rossby number characteristic of Saturn.  All of our simulations show a strong prograde equatorial jet outside the tangent cylinder. At low Ekman numbers several alternating jets form in each hemisphere inside the tangent cylinder. To analyse jet scaling of our numerical models and of Jupiter and Saturn, we extend Rhines scaling based on a topographic $\\beta$-parameter, which was previously applied to an incompressible fluid in a spherical shell, to compressible fluids. The jet-widths predicted by this modified Rhines length are found to be in relatively good agreement with our numerical model results and with cloud tracking observations of Jupiter and Saturn. ",
        "introduction": "The surface flows of the gas giants Jupiter and Saturn are dominated by strong zonal motions (i.e. azimuthal flows). Zonal wind profiles at the surface are obtained by tracking cloud features with ground-based and space observations \\citep[e.g.][]{Sanchez00,Porco03,Porco05,Vasavada05}.  As shown in Fig.~\\ref{fig:vpPlanets}, these zonal winds form a differential rotation profile with alternating eastward and westward flows. Both gas giants feature a strong prograde equatorial jet flanked by several weaker alternating secondary jets. Jupiter's central equatorial jet reaches a maximum velocity around  100-140 m/s and covers latitudes within $\\pm 15^\\circ$. However the region of fast equatorial zonal flow extends to roughly $\\pm 25^\\circ$, and features several strong prograde and retrograde jets that display marked equatorial asymmetry. The secondary undulating zonal winds at higher latitudes are weaker ($\\sim 10-20$ m/s) and narrower than the equatorial flow. Saturn's equatorial flow is stronger than that of Jupiter, wider, and more symmetrical, with a single equatorial jet extending roughly $\\pm 30^\\circ$ in latitude. It is flanked by several secondary jets in each hemisphere. These winds are significantly shifted towards the prograde direction when assuming the rotation period measured by Voyager \\citep{Desch81}. However, the suitability of this so-called System III rotation period, which is based on Saturn's kilometric radio (SKR) emissions, for characterizing the mean planetary rotation rate has been questioned. The current Cassini space mission has measured an apparent 6 minute increase in the SKR rotation period since Voyager \\citep[e.g.][]{Sanchez05}. This simple change of 1\\% in the rotation rate can substantially modify the zonal wind speed (typically around 20\\% near the equator). Since such changes in the planetary rotation rate are unlikely \\citep{Heimpel12}, alternative rotation systems have been proposed. Here we adopt the System IIIw rotation period of  \\cite{Read09}, which is based on an analysis of potential vorticity. This System IIIw rotation period yields more symmetric zonal flows between the prograde and the retrograde directions (see Fig.~\\ref{fig:vpPlanets}). The amplitude of the equatorial flow then reduces from 450 m/s in the System III rotation period to 370 m/s in System IIIw.   \\begin{figure}[t] \\centering \\includegraphics[width=8.8cm]{Fig1.pdf} \\caption{Surface zonal flow profiles for Jupiter (left panel) and Saturn (right panel). Jupiter's zonal flow profiles come from Cassini's data by \\cite{Porco03} and \\cite{Vasavada05}. Saturn's profiles in the two rotations systems have been derived by \\cite{Read09}. The conversion between zonal winds velocities  in m/s and Rossby numbers is given by $u/\\Omega r_o$, $r_o$ being the planetary radius at the $1$ bar level.} \\label{fig:vpPlanets} \\end{figure}  Modelling of zonal wind dynamics can be categorised in two main approaches. In ``shallow models'', which are typically based on the hydrostatic approximation of fluid dynamics, the dynamics is confined to a very thin layer close to the cloud level \\citep[e.g.][]{Vallis06, Liu11}. In this approach, zonal winds are maintained by turbulent motions coming from several possible physical forcings that occur at the stably-stratified cloud level (e.g. latent heat release or solar radiation). These shallow models reproduce several features observed on Jupiter and Saturn and notably the alternating direction of the zonal winds \\citep{Williams78,Cho96}. While earlier models yielded a retrograde equatorial jet for both gas giants, prograde equatorial flow has been obtained in more recent models via the inclusion of additional forcing mechanisms, such as water vapor condensation \\citep{Lian10} or enhanced radiative damping \\citep[e.g.][]{Liu11}.  In an alternative approach, zonal winds may be driven by deep-seated convection. This is supported by the latitudinal thermal emission profiles of Jupiter and Saturn, which are relatively flat, implying that deep-seated, outward directed heat flow exceeds absorbed solar heat for both planets \\citep{Ingersoll76, Pirraglia84}. Although deep convection must drive the dynamos that produce the global magnetic fields of Jupiter and Saturn, it is thought that fast zonal flows would be strongly attenuated at depth by the magnetic braking resulting from the strong increase in the electrical conductivity \\citep{Guillot04,French12}. Based on an estimate of the associated Ohmic dissipation, \\cite{Liu08} concluded that the zonal flows must therefore be confined to a relatively thin layer (i.e. $0.96\\, R_J$ and $0.85\\,R_S$). This magnetic exclusion of fast zonal flow to an outer region has been demonstrated in recent dynamo models that include radially variable electrical conductivity \\citep{Heimpel11, Duarte13}.  The deep convection hypothesis has been tested by 3-D numerical models of turbulent convection in rapidly-rotating spherical shells \\citep[e.g.][]{Christensen01, Heimpel05}. Rapid rotation causes convection to develop as axially-oriented quasi-geostrophic columns \\citep{Busse70}. This columnar flow gives rise to Reynolds stresses, a statistical correlation between the convective flow components that feeds energy into zonal flows \\citep[e.g.][]{Busse94}. The usual prograde tilt of the convective columns goes along with a positive flux of angular momentum away from the rotation axis that yields an eastward equatorial flow \\citep[e.g.][]{Zhang92}. While such models can easily reproduce the correct direction and amplitude of the equatorial jets observed on Jupiter and Saturn, they usually fail to produce multiple high-latitude jets \\citep{Christensen02}. Numerical models in relatively deep layers and moderately small Ekman numbers (i.e. aspect ratio $r_i/r_o=0.6$ and $E=10^{-4}-10^{-5}$) typically produce only a pair of jets in each hemisphere \\citep{Christensen01,Jones09,Gastine12}.  Reaching quasi-geostrophic turbulence in a 3-D model of rotating convection is indeed numerically very demanding. These numerical difficulties can be significantly reduced by using the quasi-geostrophic approximation which includes the effects of the curvature of the spherical shell (i.e. the topographic $\\beta$-effect) while solving the flow only in the 2-D equatorial plane \\citep{Aubert03,Schaeffer05NP}. For instance, the two-dimensional annulus model of rotating convection considered by \\cite{Jones03,Rotvig06} and \\cite{Teed12} allows very low Ekman numbers to be reached. Multiple alternating zonal flows have been found and the typical width of each jet scales with the Rhines length \\citep[e.g.][]{Rhines75,Gurarie02,Read04,Suko07}. Independently of these models, the 3-D simulations of \\cite{Heimpel05} and \\cite{Heimpel07} were computed with lower Ekman numbers and larger aspect ratio ($r_i/r_o \\ge 0.85$)  than previous numerical models \\citep[e.g.][]{Christensen01}. These simulations show clear evidence of multiple jets inside the tangent cylinder. While these simulations produce fewer bands than observed in the  quasi-geostrophic simulations, at least in part due to the unavoidable computational limitations, the width of each zonal band follows the Rhines scaling \\citep[][hereafter HA07]{Heimpel07}.  Most of the previous models have employed the Boussinesq approximation where compressibility effects are simply ignored. In giant planets, however, the density \\modif{increases by more than 1000 in the molecular envelope} \\citep{Guillot99,French12} and the applicability of the topographic Rhines scaling is therefore questionable \\citep{Evonuk08}. More recent anelastic models that allow to incorporate compressibility effects have nevertheless shown several similarities with the previous Boussinesq studies. While the density contrast affects the small-scale convective motions, the large-scale zonal flows are relatively unaffected by the density stratification \\citep{Jones09,Gastine12}. However, these results seem incompatible with the new vorticity source introduced by compressibility, which may change the definition of the $\\beta$-effect \\citep[e.g.][]{Ingersoll82,Glatz09,Verhoeven14}. As pointed out by \\cite{Jones09}, defining a consistent Rhines scaling in a compressible fluid is not straightforward.  To tackle this problem, we consider numerical models of rapidly-rotating convection in a compressible fluid. We present numerical simulations of a thin spherical shell ($r_i/r_o=0.85$) with a strong density stratification ($\\rho_\\text{bot} /\\rho_\\text{top} \\simeq 150$). A series of simulations, with decreasing Ekman number, are employed to gradually reach the multiple-jets regime. To analyse the zonal flow scaling, we present a consistent derivation of Rhines scaling in compressible fluids.  The paper is organised as follows. In section~\\ref{sec:model}, we present the anelastic formulation and the numerical method. The results of the numerical simulations are presented in section~\\ref{sec:results}. A new derivation of compressible Rhines scaling is presented in section~\\ref{sec:rhines}. In section~\\ref{sec:interiormodels}, we concentrate on the applications to numerical models and gas giants, before concluding in section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have investigated several numerical simulations of rapidly-rotating convection in spherical shells. We have considered strongly stratified models with $\\rho_\\text{bot}/\\rho_\\text{top} \\simeq 150$ in spherical shells with a large aspect ratio of $r_i/r_o=0.85$. Extending the previous Boussinesq study by HA07, we have explored a range of Ekman numbers ($E=3\\times 10^{-4}-3\\times 10^{-6}$) to allow the formation of multiple high latitude jets.  At the lowest Ekman number considered here, a strong prograde equatorial zonal flow is flanked by several alternating jets in each hemisphere. This confirms previous findings by \\cite{Jones09} that compressible models can also maintain multiple zonal bands, provided low Ekman numbers and thin spherical shells are considered. The number of jets in our simulations is far fewer than observed for Jupiter. Smaller Ekman numbers, and even thinner shells would result in a greater number of narrower jets. However, both of these properties require computational expense that is presently impractical. Following our previous parameter studies of compressible convection \\citep{Gastine12,Gastine13}, the zonal flow profiles obtained here show many similarities with their Boussinesq counterparts \\citep[e.g.][HA07]{Heimpel05}. In agreement with the theory of geostrophic turbulence, the jet widths of these previous Boussinesq models have been found to scale with the Rhines length, i.e. $k_\\beta^{-1} = \\sqrt{2U/\\beta}$, $\\beta$ being the relative variation of the height of the spherical shell with the distance to the rotation axis \\citep{Rhines75}. The similarities in the zonal flow profiles of Boussinesq and anelastic models would therefore suggest that this topographic Rhines scale could be also applicable to compressible convection. This seems however at odds with the two-dimensional anelastic models by \\cite{Evonuk08} and \\cite{Glatz09} which suggested a possible modification of the Rhines lengthscale due to compressibility effects.  To address this issue, we have derived here the expression of the $\\beta$-effect in a 3-D spherical shell filled with a compressible fluid. Confirming the calculations by \\cite{Ingersoll82}, the compressional $ \\beta$-effect (i.e. $\\beta_\\rho$) is found to be proportional to the relative variation of the integrated mass with the distance to the rotation axis. Similarly to the classical topographic Rhines scaling (i.e. $\\beta_h$), the compressional $\\beta$-effect marks a sign reversal at the tangent cylinder. In the thin spherical shells considered here, the $\\beta_\\rho$ and $\\beta_h$ profiles are found to be very similar inside the tangent cylinder. This explains the similarities of the high-latitude jets observed in Boussinesq and anelastic models. The main differences between topographic and compressional $\\beta$-effects occur outside the tangent cylinder where the density gradient is mainly perpendicular to the rotation axis. However the Rhines scaling does not seem to be applicable there as there is only one single equatorial jet which seems to be always controlled by the geometry of the spherical shell (see also HA07). An efficient test of the compressional Rhines scale for $s> r_i$ would therefore require multiple zonal bands outside the tangent cylinder. While those low-latitude jets are frequently observed in the transient phases of the 3-D numerical models, they always merge on longer time resulting in one single prograde equatorial flow flanked by a retrograde jet attached to the tangent cylinder \\citep[e.g.][]{Sun95,Christensen01}. Although numerically challenging,  3-D models in thicker spherical shells (i.e. $r_i/r_o \\leq 0.6$) at low Ekman numbers (i.e. $E \\sim 10^{-6}$) may help to check the stability of the zonal winds outside the tangent cylinder. As the differences between $\\beta_h$ and $\\beta_\\rho$ are also more pronounced in thicker shells, it would be interesting to check the applicability of the Rhines scaling to these low-latitude jets.  The compressional $\\beta$-effect has also been applied to Jupiter's and Saturn's surface winds. For Jupiter, $\\beta_h$ and $\\beta_\\rho$ are relatively similar and both predict jet-widths in relatively good agreement with the observed values. For Saturn, the misfit between predictions and observations is on average larger. The compressional Rhines scale gives nevertheless a better estimate of the jet-width than the topographic one. Excluding the equatorial jet, the best fit is achieved for a range of radius ratios of $\\eta\\simeq 0.85-0.9$ for Jupiter and $\\eta\\simeq 0.55$ for Saturn. These values are on the low-side compared to the estimates derived by \\cite{Liu08}. However these depths are relatively poorly constrained by the Rhines scaling owing to the small changes when the radius ratio is varied. In addition, the factor 2 entering the Rhines scale definition is not firmly established as it comes from an \\textit{ad-hoc} averaging of the orientation of the Rossby waves \\citep{Rhines75}. \\cite{Suko07} therefore argues that this factor should be replaced by a rather elusive coefficient of proportionality. It is therefore difficult to put a strong constraint on the zonal flow depth based on $\\beta_\\rho$ only. \\modif{Using the compressional Rhines scale also allows to construct an asymptotic scaling for the number of zonal jets. However, estimating this number requires to employ an asymptotic scaling for the zonal flow amplitude. The disagreement between the different buoyancy power scalings published to date suggests that the zonal flow might obey a more complex scaling. Thus, we propose a simple best-fit scaling law that successfully predicts the correct zonal flow amplitude and number of bands for both the numerical models and the gas giant planets. Although encouraging, the predicted dependence on the control parameters $Ra^*$, $E$ and $Pr$ is still tentative and needs further theoretical investigations to confirm and justify these exponents.}  \\modif{The exact structure of the zonal jets in the gas giant interiors is still a matter of debate. The anelastic models by \\cite{Kaspi09} for instance suggest an important variation of the zonal flow amplitude along the direction of the rotation axis that is not captured by our quasi-geostrophic zonal flow profiles. Testing the applicability of the compressional Rhines scaling to this kind of baroclinic jets is therefore an interesting prospect that would need to be addressed by future models.}  Finally a deeper understanding of the jet scaling would also require the study of integrated models that include the deep-interior dynamics. Recent Boussinesq and anelastic dynamo models with radially decreasing electrical conductivity have shown that the penetration of the zonal flows is limited by Lorentz forces that increase strongly at depth \\citep{Stanley09,Heimpel11,Duarte13}. While those models still show a strong prograde equatorial zonal flow, they lack well-defined high-latitude jets. Further investigations are therefore needed to analyse the depth and the stability of these high-latitude zonal winds."
    },
    "1402/1402.4090_arXiv.txt": {
        "abstract": "We construct empirical models of the diffuse gamma-ray background toward the Galactic Center. Including all known point sources and a template of emission associated with interactions of cosmic rays with molecular gas, we show that the extended emission observed previously in the Fermi Large Area Telescope data toward the Galactic Center is detected at high significance for all permutations of the diffuse model components. However, we find that the fluxes and spectra of the sources in our model change significantly depending on the background model. In particular, the spectrum of the central Sgr A$^\\ast$ source is less steep than in previous works and the recovered spectrum of the extended emission has large systematic uncertainties, especially at lower energies. If the extended emission is interpreted to be due to dark matter annihilation, we find annihilation into pure $b$-quark and $\\tau$-lepton channels to be statistically equivalent goodness of fits. In the case of the pure $b$-quark channel, we find a dark matter mass of $39.4\\left(^{+3.7}_{-2.9}\\rm\\ stat.\\right)\\left(\\pm 7.9\\rm\\ sys.\\right)\\rm\\ GeV$, while a pure $\\tau^{+} \\tau^{-}$-channel case has an estimated dark matter mass of $9.43\\left(^{+0.63}_{-0.52}\\rm\\ stat.\\right)(\\pm 1.2\\rm\\ sys.)\\ GeV$. Alternatively, if the extended emission is interpreted to be astrophysical in origin such as due to unresolved millisecond pulsars, we obtain strong bounds on dark matter annihilation, although systematic uncertainties due to the dependence on the background models are significant. ",
        "introduction": "The Milky Way's Galactic Center (GC) harbors an extremely dense astrophysical environment, with thousands of high-energy sources detected in the X-ray within the inner 0.3$^\\circ$~\\cite{Muno:2008qy}, as well as numerous gamma-ray emitting point sources~\\cite{Nolan2012}. In addition, the GC is expected to harbor high densities of dark matter (DM) with a power-law increase in density toward the center, leading it to be among the best places in which to find signals of DM particle annihilation or decay \\cite{Springel:2008by}. A leading candidate for cosmological dark matter is a thermally produced weakly interacting massive particle (WIMP) that can arise in many extensions of the Standard Model of particle physics, whose annihilation is related to their production in the early Universe~\\cite{Feng:2010gw}.  Several groups have found strong evidence for extended emission in the gamma ray from the GC using data from the Large Area Telescope (LAT) aboard the {\\it Fermi Gamma-ray Space Telescope}. It has been shown that the extended emission is consistent with the spatial profile expected in DM halo formation simulations, the flux is consistent with the annihilation rate of thermally produced WIMP DM, and the spectrum is consistent with the gamma rays produced in the annihilation of $\\sim\\!\\! 10 - 30$ GeV DM to quarks or leptons~\\cite{Goodenough:2009gk,Hooper:2010mq,Hooper:2011ti,Boyarsky:2010dr,Abazajian:2012pn,Gordon:2013vta,Macias:2013vya}. This triple consistency of the gamma-ray extended-source signal in the GC with the WIMP paradigm has generated significant interest. In addition, there are claims of signals consistent with the DM origin interpretation in the ``inner Galaxy'' \\cite{Hooper:2013rwa}, and in unassociated point sources \\cite{Berlin:2013dva}. The required dark matter mass and annihilation cross section is consistent with the constraints from Milky Way dwarf galaxies~\\cite{Ackermann:2013yva}.  \\begin{figure*}[h!t!] \\begin{sideways} \\makebox[1.6truein][c]{$\\quad$ Diffuse Sources' Residual} \\end{sideways} \\includegraphics[width=1.68truein]{Diffuse_resid_07.pdf} \\includegraphics[width=1.68truein]{Diffuse_resid_09.pdf} \\includegraphics[width=1.68truein]{Diffuse_resid_10.pdf} \\includegraphics[width=1.68truein]{Diffuse_resid_11.pdf}\\\\ \\begin{sideways} \\makebox[1.6truein][c]{$\\quad$ Diffuse Models} \\end{sideways} \\includegraphics[width=1.68truein]{Diffuse_model_07.pdf} \\includegraphics[width=1.68truein]{Diffuse_model_09.pdf} \\includegraphics[width=1.68truein]{Diffuse_model_10.pdf} \\includegraphics[width=1.68truein]{Diffuse_model_11.pdf}\\\\ \\begin{sideways} \\makebox[2.1truein][c]{\\ $\\qquad$\\ \\ GCE Source Residual} \\end{sideways} \\includegraphics[width=1.68truein]{MSP_resid_07.pdf} \\includegraphics[width=1.68truein]{MSP_resid_09.pdf} \\includegraphics[width=1.68truein]{MSP_resid_10.pdf} \\includegraphics[width=1.68truein]{MSP_resid_11.pdf}  \\caption{For the full model, 2FGL+2PS+I+MG+ND+GCE (see text and Table \\ref{logliketable}), we show here the multicomponent diffuse model (the combined I+MG+ND) residuals, i.e., the counts subtracting all model components other than the I+MG+ND components (top row), the multicomponent diffuse model, I+MG+ND, (middle row), and the GCE source residuals within our ROI (bottom row), at $|b| < 3.5^\\circ$ (vertical axis) and $|\\ell| < 3.5^\\circ$ (horizontal axis). The maps are shown on the same color scale to show the components' relative strength for the counts per pixel, Gaussian filtered spatially with $\\sigma=0.3^\\circ$. \\label{diffusemap}} \\end{figure*}   Alternatively, the high density of compact objects, cosmic-ray emission, and other astrophysical activity in the GC is also expected to be a considerable source of gamma-ray emission. The massive GC Central Stellar Cluster may harbor a significant millisecond pulsar (MSP) population that can have similar gamma-ray flux and spectrum as the observed extended source in the GC \\cite{Abazajian:2010zy}.  There is also a significant detection of gamma-ray emission associated with molecular gas as mapped by the 20 cm radio map toward the GC \\cite{YusefZadeh:2012nh}. In Ref.~\\cite{YusefZadeh:2012nh}, the 20 cm map had the strongest statistical detection of the diffuse source templates studied, and is interpreted as bremsstrahlung emission of high-energy electrons interacting with the molecular gas (MG). In addition, the gamma-ray point source associated with Sgr A$^\\ast$ is among the brightest sources in the gamma-ray sky. Sgr A$^\\ast$'s spectrum from low- to high-energy gamma rays has been modeled to originate from cosmic-ray protons transitioning from diffusive propagation at low energies to rectilinear propagation at high energies \\cite{Chernyakova:2011zz,Linden:2012iv}. Interestingly, the potential confusion between pion decay, pulsar spectra and DM annihilation was studied well before the launch of the Fermi LAT \\cite{Baltz:2006sv}.    In this paper, we perform a detailed analysis of the nature of the extended gamma-ray source from the GC region, which we designate as the GC extended (GCE) source, the point sources in the GC, as well as the diffuse emission associated with the 20 cm MG map. We focus on a region of interest (ROI) of $7^\\circ\\times 7^\\circ$ centered at the GC. Since there have been detections of all of these sources independently and their spatial information overlaps, we perform a new analysis which consistently includes all of these sources---extended, point-like, and diffuse---as well as their uncertainties determined by the data.  In addition, including systematic and statistical uncertainties, we determine the best fit particle masses and annihilation channels if the GCE is interpreted as DM. Conversely, in the case of interpreting the GCE source as an MSP population, we discuss the number of MSPs required within our ROI and we also place strong limits on DM annihilation cross sections. ",
        "conclusions": "\\label{conclusions}  We have presented the results of a large set of analyses of the nature of point source, diffuse and extended source gamma-ray emission toward the Milky Way's Galactic Center as observed by the Fermi LAT. We have included all known point sources toward the GC as well as a template of the molecular gas based on radio emission. In all cases, we find a highly statistically significant robust detection of an extended source consistent with dark matter annihilation and/or a population of millisecond pulsars in the GC. However, the detailed spectrum of this extended source depends strongly on the background (diffuse source) models.  The spectrum of the source associated with Sgr A$^\\ast$ is less steep than in previous work, owing to the new extended and diffuse sources. In the case of a dark matter annihilation interpretation of the GC extended source, the particle mass is very precisely determined given an annihilation channel, though systematic uncertainties in the diffuse emission introduce significant systematic uncertainties.  The $b$-quark or $\\tau$-lepton channels are almost equally preferred, but with different particle masses.  For annihilation into $b$ quarks we find $m_\\chi=39.4\\left(^{+3.7}_{-2.9}\\rm\\ stat.\\right)\\left(\\pm 7.9\\rm\\ sys.\\right)\\rm\\ GeV$, $\\langle \\sigma v\\rangle_{bb} = (5.1\\pm 2.4)\\times 10^{-26}\\rm\\ cm^3\\ s^{-1}$. For the $\\tau^{+} \\tau^{-}$ channel we find $m_\\chi = 9.43\\left(^{+0.63}_{-0.52}\\rm\\ stat.\\right)(\\pm 1.2\\rm\\ sys.)\\ GeV$, $\\langle \\sigma v\\rangle_{\\tau^{+} \\tau^{-}} = (0.51\\pm 0.24)\\times 10^{-26}\\rm\\ cm^3\\ s^{-1}$. These annihilation rates are lower than, but close to the annihilation rates that are excluded by combined dwarf galaxy analyses~\\cite{Ackermann:2013yva} and collider searches~\\cite{ATLAS:2012ky}. Future combined dwarf galaxy analyses may be sensitive to this parameter space~\\cite{KoushiappasAspen2013,DrlicaWagner:2013,He:2013jza}.  Once confirmed, measurements of the isotropic extragalactic background can yield further information on, e.g., the smallest halo mass \\cite{Ng:2013xha}.  It has been pointed out that bremsstrahlung will modify the gamma-ray spectra appreciably \\cite{Cirelli:2013mqa}, and our tests find that they increase the inferred particle masses in the $b \\overline{b}$ or $\\tau^{+} \\tau^{-}$ channels.  While the extended source is robustly detected, we caution that the shape of the rise and fall of the spectrum ($E^2dN/dE$), as shown in Figs.~\\ref{GCEspectrum} and \\ref{systematicsspectrum}, is highly model dependent.  When interpreting all of the GCE emission as astrophysical, we find stringent limits on dark matter annihilation, but they are highly model dependent. In this sense, the combined dwarf limits are still the most robust.  To explain the diffuse GCE emission with unresolved MSPs, we estimated (using the gamma rays from 47 Tuc as a reference) that there need to be about 3000 to 5000 MSPs within the ROI (1 kpc by 1 kpc box towards the GC). This is a large number compared to the typical number of MSPs in globular clusters but the total stellar content is also much larger in this region. We have also highlighted the possibility that multiple sources may contribute to the GCE.  While we have characterized some of the systematic uncertainty associated with modeling of the diffuse background, we emphasize that our treatment is far from exhaustive. Further multiwavelength study of the Milky Way's Galactic Center is essential to understanding the nature of the numerous sources in this highly dense astrophysical region. Even so, the detection of the GCE source is fairly robust to differences in the background modeling, and though the extended emission in gamma rays studied here is consistent with a pure astrophysics interpretation, the extended emission's consistency in morphology, spectrum and flux with a dark matter annihilation interpretation remains extremely intriguing.   \\appendix*"
    },
    "1402/1402.3509_arXiv.txt": {
        "abstract": "Active Galactic Nuclei (AGN) are axisymmetric systems to first order; their observed properties are likely strong functions of inclination with respect to our line of sight, yet the specific inclinations of all but a few AGN are generally unknown. By determining the inclinations and geometries of nearby Seyfert galaxies using the kinematics of their narrow-line regions (NLRs), and comparing them with observed properties, we find strong correlations between inclination and total hydrogen column density, infrared color, and H$\\beta$ full-width at half maximum (FWHM). These correlations provide evidence that the orientation of AGN with respect to our line of sight affects how we perceive them, beyond the Seyfert 1/2 dichotomy. They can also be used to constrain 3D models of AGN components such as the broad-line region and torus. Additionally, we find weak correlations between AGN luminosity and several modeled NLR parameters, which suggests that the NLR geometry and kinematics are dependent to some degree on the AGN's radiation field. ",
        "introduction": "\\label{sec1}   Seyfert galaxies represent a relatively moderate luminosity ($L_{bol} \\approx 10^{43} - 10^{45}$ erg s$^{-1}$), nearby (z $\\leq$ 0.1) subset of the overall collection of AGN, which exhibit a dichotomy of broadened and unbroadened permitted emission lines. This dichotomy has led the Seyfert class of AGN to be divided into two groups \\citep{Kha74}, with Seyfert 1s exhibiting spectra containing broad (full width at half-maximum [FWHM] $\\geq$ 1000 km s$^{-1}$) permitted lines, narrower (FWHM $\\leq$ 1000 km s$^{-1}$) forbidden lines, and distinct, non-stellar optical and UV continua, and Seyfert 2s containing only narrow permitted and forbidden emission lines, with their optical and UV continua being dominated by the host galaxy.  \\citet{Ost78} suggested that Seyfert 1s and 2s were physically the same objects, but that the BLRs of Seyfert 2s are obscured along our line of sight (LOS). The strongest observational evidence for a unified model stems from work by \\citet{Ant85}, who discovered that weak, polarized broad emission lines do exist in Seyfert 2 galaxies by analyzing the polarized spectrum of the Seyfert 2 NGC 1068. This observation led to the generation of a unified AGN model, with AGN being similar objects, obscured by a toroidal structure of gas and dust in Seyfert 2s \\citep{Ant93}, viewed from different angles. Observing an AGN from near the axis of the torus would allow for the viewing of the central source, the BLR, and the NLR, and produce the spectrum of a Seyfert 1, while observing the same AGN from near the plane of the torus would leave only the NLR and a small amount of emission from the nuclear source visible, producing a Seyfert 2 spectrum. It is important to note that this is an oversimplified scenario, in which each AGN with a given luminosity is contained within an identical, smooth-density torus. In reality, the torus around each AGN likely varies in covering factor, or how much sky at the AGN center is covered by obscuring material, and has a clumpy composition that does not create a hard cut-off between observing a Type 1 or Type 2 AGN, but instead only reduces the probability to directly view the AGN as our line of sight strays further from the torus axis \\citep{Eli12}. The NLR is then the radiation that has escaped the torus, forming a biconical structure likely governed by the torus covering factor, with an axis that is perpendicular to the plane of the torus.  In our previous paper (\\citealt{Fis13a}, henceforth Paper I), we measured radial velocities of [O~III] emission knots across NLRs for 47 Seyfert AGN observed with {\\it Hubble Space Telescope} ({\\it HST}) STIS G430L/M gratings. From our measurements, 12 AGN showed clear evidence for  kinematics dominated by biconical outflow, where knots accelerated out from the inner nucleus, reached a terminal velocity, and then decelerated back toward the AGN systemic velocity. In order to determine the geometry, including the inclination of their NLR outflows, and thus the inclination of their torus, (both of which generally unknown for all but a few AGN) a kinematic model was fit to the observed NLR kinematics using a simple, linear velocity law. Using available imaging and kinematic data for each AGN to create an initial geometry, we refined the model in an iterative process until a best fit to the kinematic data was obtained. Including 5 AGN previously modeled by our group (\\citealt{Cre00a,Cre00b,Rui01,Das05,Das06,Cre10a,Fis10,Fis11}), we now have an initial sample of 17 AGN inclinations, from which we can begin to determine how their observed properties vary as a function of polar angle (i.e. the inclination of the torus) with respect to the accretion disk and/or torus axes, and identify correlations that probe the structure of the AGN components by comparing our kinematic model results with observed multiwavelength properties of these AGN.  Section 2 describes the parameters and how they were measured. Section 3 illustrates all relevant comparisons between modeled and observed AGN parameters. Sections 4 notes significant non-correlations. Section 5 contains discussion and final conclusions. ",
        "conclusions": "\\label{sec4}  With inclinations and geometries of 17 Seyfert galaxies, our model results suggest that Seyfert 1 AGN are inclined further toward our LOS than Seyfert 2 AGN. Knowing the inclinations of these AGN allows us to determine how their observed properties vary as a function of polar angle with respect to the accretion disk and/or torus axes. We have established a strong connection between polar angle and column density, which probes the structure of the obscuring torus, the ionized gas outflows, and the connection between the two; polar angle and broad H$\\beta$ FWHM, suggesting a non-spherical component to the BLR kinematics; and polar angle and mid-IR color (5 $\\mu$m/30 $\\mu$m flux), presumably indicating a better view of the hot inner ``throat'' of the obscuring torus in Seyfert 1s compared to Seyfert 2s.  With the discovery of three independent correlations on inclination, we can see that our line of sight affects how we observe specific parameters of our sample and thus that the obstructing torus surrounding the central engine plays a significant role in how we observe these AGN. Additionally, we find correlations between AGN luminosity and the NLR kinematics and geometry, but not between opening angle, molecular hydrogen mass surrounding the torus and the NLR kinematics and geometry.  Our inclination results agree with and help quantify the unified model, as $N_H$, IR color, and $H\\beta_{FWHM}$ measurements each show a strong correlation with viewing angle. These results can thus provide valuable constraints on specific torus models (e.g. \\citealt{Fri06,Ram09,Rot12,Sta12}) and BLR models (e.g. \\citealt{Col06,Den12,Pan12}).  Establishing accurate, continuous correlations between AGN inclination and independent, observable parameters may allow us to use these comparatively simple to obtain observational measurements in the future as a proxy for inclination in AGN with unmodelable kinematics. What other observable parameters could be related to inclination? Could the amount of polarized light depend on inclination? As \\citet{Ant93} observed, polarization of type 1 galaxies tends to be parallel to the AGN jet axis and polarization of type 2 galaxies tends to be perpendicular to the jet axis. Could this also apply to the torus axis, such that a correlation between increasing perpendicular polarization and polar angle exists?  One parameter that certainly requires further investigation is the role of the host galaxy in obscuring the AGN. As mentioned earlier, the maximum sample size is one less than our total AGN sample size, as we omit NGC 5506\\footnote{Paper I stated that the classification of NGC 5506 was contested, with labels of Narrow-Line Seyfert 1 (NLS1;\\citealt{Nag02}) / Sy 1.9 \\citep{Mai95} / Sy 2 AGN \\citep{Tri10} all being used to describe the AGN. However, \\citet{Tri10} notes that the alternate labels do not correspond to what is seen optically in NGC 5506 and we have removed the other labels from contention in our analysis.} from our comparisons because it contains a highly inclined (76$^{\\circ}$) host disk. Imaging depicts the NLR as a single cone near perpendicular to the host disk, and in Paper I, we modeled the polar angle of the NLR to be 80$^{\\circ}$ from LOS. However, it is unknown how much of the visible NLR is extinguished due to the host disk orientation, although X-ray observations of NGC 5506 have shown that at least a portion of the NLR is extinguished as X-ray emission from a second NLR cone can be clearly detected on the opposite side of the nucleus from the visible NLR \\citep{Zen09}. As we found no correlation between NLR and host disk orientation in Paper I, it is possible that the near edge-on host disk would play a major role in obscuring portions of a moderately inclined NLR, and the interpretation of only the unobscured kinematics would thus affect the resultant kinematic model. If the NLR were to be inclined only 40 degrees from our LOS, per Fe K$\\alpha$ analysis in \\citet{Gua10}, then the NLR may be opened wide enough that a portion of the NLR intersects with the disk. Several other modeled AGN experience intersections between the host and NLR (i.e. Mrk 3, Mrk 34, Mrk 573) which illustrate that the NLR is radiation bound in an intersection scenario and emission near this intersection does not continue to radially progress through the host disk. As such, an observer near the edge of the host disk of each of these AGN would likely be unable to observe an integral portion of the NLR kinematics. Additionally, while host disks likely play a role in the amount of NLR emission that is extinguished in each AGN, they do not contain hydrogen column densities significant enough to account for the X-ray/UV absorption observed over each AGN. Comparing host disk inclination to hydrogen column density in our sample results in a correlation coefficient of $r = 0.034$ and a corresponding p-value of $P_{c}(r,N) = 89\\%$. Thus, at its current NLR orientation, NGC 5506 has a hydrogen column density orders of magnitude lower than all other AGN at high inclinations, implying that it should have an inclination closer to our LOS. A polar angle of $\\theta = 40^{\\circ}$ would place the $N_{H} = 2 \\times 10^{22}$ column density measurement much closer to the given correlation.  Within our limited range of luminosities, we find that outflow velocity, NLR height, and turnover radius each appear to be influenced by the luminosity of the AGN. Additionally, NLR height and turnover radius are well correlated, which means whatever is determining the ultimate reach of photons from the AGN also has some role in the dynamics of the system, extending the acceleration of NLR clouds to further distances. Could these relations provide insight on the origin of the NLR? If the NLR emitting gas primarily originates from a stream of ionizing radiation unleashed upon the galaxy, ionizing and accelerating ambient material that it encounters, the one property that should affect the geometry of the NLR the greatest would be AGN luminosity. To more properly understand this relation requires additional models to further constrain the strength of the given correlations and further study of the orientation between the NLR radiation and the plane of the host galaxy. AGN that display such an intersection (i.e. Mrk 78, NGC 4151) typically result in a more extended NLR and are more likely to host in situ acceleration of NLR gas \\citep{Fis10}.  From our relation between bolometric luminosity and maximum outflow velocity, we see that outflow kinematics possibly cease at L$_{bol} < 10^{41}$ erg s$^{-1}$. {\\it HST} STIS H$\\alpha$ studies similar to our own in \\citet{Wal08} depict observed low-ionization nuclear emission-line regions (LINERs), AGN with luminosities generally lower than Seyferts, containing possible rotation (or 'ambiguous'; see Paper I) kinematics with blueshifted velocities $<$ 250 km s$^{-1}$. Though these velocities may be due to projected outflows, they could instead be a combination of rotation and ionization via star formation or post-AGB stars \\citep{Yan12}.  Finally, we note that although the correlations that we found are intriguing, they are based on a small number of data points. More confirmations of radial outflows in AGN and corresponding kinematic models would result in more inclinations for testing these correlations. More observations of Seyfert 2s at higher X-ray energies would help replace lower limits on $N_H$ with actual values and additional X-ray observations of our current sample over several epochs would allow us to determine if and how much $N_H$ varies over time."
    },
    "1402/1402.3201_arXiv.txt": {
        "abstract": "We consider a family of cosmological models in which all mass is confined to a regular lattice of identical black holes. By exploiting the reflection symmetry about planes that bisect these lattices into identical halves, we are able to consider the evolution of a number of geometrically distinguished surfaces that exist within each of them. {\\edit We find that the evolution equations for the reflection symmetric surfaces can be written as a simple set of Friedmann-like equations, with source terms that behave like a set of interacting effective fluids.} We {\\edit then} show that gravitational waves are effectively trapped within small chambers for all time, and are not free to propagate throughout the space-time. Each chamber therefore evolves as if it were in isolation from the rest of the universe. We call this phenomenon ``piecewise silence''. ",
        "introduction": "\\label{sec:int}  Since its birth, classical general relativity has been applied to the construction of cosmological models. This is usually done by specifying the symmetries that are expected to exist on large scales, and then looking for solutions of Einstein's field equations that exhibit those symmetries. The Cosmological Principle, that the Universe is spatially homogeneous and isotropic on large scales, then leads one to the Friedmann-Lema\\^itre-Robertson-Walker (FLRW) solutions. Taken as a model for the Universe, the FLRW solutions of Einstein's equations have been shown to be remarkably consistent with a wide array of observations, ranging from the cosmic microwave background \\cite{planck}, all the way through to the Hubble diagrams constructed using nearby supernovae \\cite{perlmutter,riess}. This consistency, however, is only allowed at the expense of including new and exotic ``dark'' components that must dominate over all other forms of matter in the Universe.  On the other hand, we know that the visible matter in the late Universe is highly inhomogeneous on small scales, being condensed largely into stars that only occupy about $10^{-30}$ of the total volume of space\\footnote{Based on a typical stellar density of $\\sim 1$g/cm$^{3}$, a baryon fraction of $\\sim 0.05$, and a critical density of $\\sim 10^{-29}$g/cm$^{3}$.}. The rest is an almost perfect vacuum. This strongly suggests the need to relax our assumptions about homogeneity, at least at some level. Bound up with this is the ``backreaction'' problem, that averaging and evolution do not commute in Einstein's theory \\cite{Andersson&Coley:2011}. This means that even if the Universe is close to being homogeneous and isotropic on large scales, there is no guarantee that it will evolve anything like the FLRW solutions of Einstein's equations (unless one adds some otherwise undetectable exotic matter fields, in order to force it to).  The problem that then needs to be addressed, in order to trust the results of the standard approach to cosmology, is to determine the large-scale evolution that emerges in space-times that are close to statistically homogeneous and isotropic on large scales, but very highly inhomogeneous on small scales. This problem has been approached in many different ways in the literature, including using cosmological perturbation theory \\cite{perts}, spatial averaging \\cite{space}, space-time averaging \\cite{space-time}, shortwave approximations \\cite{green}, exact solutions \\cite{bull}, and approximate solutions \\cite{larena}, to name but a few. Our approach is to develop a family of cosmological models that contains a regular array of identical black holes only. These models have the great benefit of admitting a time-symmetric initial value problem that can be solved exactly \\cite{Clifton_etal:2012}, as well as allowing for the exact evolution of some high-symmetry curves to be found \\cite{Clifton_etal:2013}.  Early work in this spirit includes that of McCrea \\cite{mccrea}, Coxeter \\& Whitrow \\cite{Coxeter&Whitrow:1950} and Lindquist \\& Wheeler \\cite{Lindquist&Wheeler:1957, geocell}. These latter authors suggested a Wigner-Seitz-type approach in which 3-dimensional space is tiled with regular polyhedra, and a mass is put at the centre of each cell. The geometry of space-time within each cell is then approximated as being Schwarzschild, and approximate junction conditions are used to obtain a large-scale evolution. The optical properties of these models have been studied in \\cite{CF1, CF2,CF3}. More recently numerical solutions to Einstein's equations have been found for these configurations in both the hyper-spherical case \\cite{bentivegna1} and the flat case \\cite{bentivegna2,yoo1,yoo2}. Even taking into account their simplified nature, these lattice models are still more realistic in many aspects than the FLRW models: Firstly, they are locally inhomogeneous in a realistic non-perturbative manner; secondly, the matter content is discrete, rather than a fluid; and thirdly, they are vacuum models, as is the real Universe at almost all points.  The focus of the present work will be on the surfaces in these models that exhibit reflection symmetry. These include the faces of the cells that constitute the original tiling, as well as a number of the planes that pass through the cell centres. While the geometry of these surfaces is a more complicated problem than those studied previously, it will also enable us to consider more general models, such as the tiling of flat space with an infinite array of cubes, as well as any other initial 2-surface that exhibits reflection symmetry, including models that are not necessarily of the regular lattice type. Our basic requirement will only be that there is an initial 2-surface with reflection symmetry. This immediately implies, in particular, that the symmetry surface itself, as well as its evolution, will be totally geodesic subspaces of the full space-time.  One feature of particular interest for relativistic cosmological models is the presence (or otherwise) of gravitational radiation. To determine whether a space-time is radiative is not straightforward in general. For our purposes, we exploit the well-known manifestly covariant and non-perturbative electromagnetic analogy for gravity \\cite{jantzen}, and say that the flux of gravitational radiation vanishes if the super-Poynting vector vanishes {\\edit \\cite{Maartens&Bassett:1998, bonilla, gomez}}. We present an argument that this condition implies that the energy carried between cells by the radiation also vanishes, despite the dimensionality of the super-Poynting vector being different from that of an energy flux. Due to the existence of a number of chambers that are entirely enclosed by these reflection symmetric surfaces, we say that our lattice models are ``piecewise silent''.  To implement this program we begin by identifying the geometric quantities that must vanish on reflection symmetric surfaces. We derive the complete (1+2)-dimensional Einstein system along their evolution. Our basic variables for this are the expansion and shear of a set of reference time-like curves, together with the non-vanishing parts of the Weyl tensor. We then evaluate the super-Poynting vector on these surfaces, and consider what this means for the propagation of gravitational waves. Unless otherwise stated, we use Greek letters $\\mu$, $\\nu$, $\\rho$ ... to denote coordinate indices, and Latin letters $a$, $b$, $c$ ... to denote tetrad indices. ",
        "conclusions": "In this paper we have studied the reflection symmetric surfaces that divide lattice universes into chambers. We have identified the restrictions that reflection symmetry imposes on kinematic and geometric quantities, and used the results to investigate the simplified equations that result for the expansion of these surfaces, as well as for the restrictions they impose on the propagation of gravitational waves.  We find that the area of the cell faces expands like the scale factor in a FLRW model that is filled with dust and radiation with energy flux, even though the space-time itself is completely devoid of any matter fields. We also find that the discrete symmetry of the configuration forces the scalar invariant $E^{ab}H_{ab}$ to vanish along all reflection symmetric surfaces, throughout the entire evolution of the model. Moreover, the only non-zero components of the magnetic Weyl tensor are shown to lie on the reflection surface itself. An explicit evaluation of the super-Poynting vector of the free gravitational field then {\\edit demonstrates} that gravitational waves are forbidden from moving between chambers, {\\edit as on the chamber boundaries there are not enough degrees of freedom in the space-time geometry to allow them to do so}. This shows that inhomogeneous discrete cosmological models with periodic boundary conditions are piecewise silent.        \\ack  We are grateful to R Tavakol for helpful discussions. TC acknowledges support from the STFC. DG is supported by the Erasmus Mundus Joint Doctorate Program by Grant Number 2011-1640 from the EACEA of the European Commission."
    },
    "1402/1402.6338_arXiv.txt": {
        "abstract": "Orphan Afterglows (OA) are slow transients produced by Gamma Ray Bursts seen off--axis that become visible on timescales of days/years at optical/NIR and radio frequencies, when the prompt emission at high energies (X and $\\gamma$ rays) has already ceased. Given the typically estimated jet opening angle of GRBs \\th$\\sim 3^{\\circ}$, for each burst pointing to the Earth there should be a factor $\\sim700$  more GRBs pointing in other directions. Despite this, no secure OAs have been detected so far. Through a population synthesis code we study the emission properties of the population of OA at radio frequencies. OAs reach their emission peak on year-timescales and they last for a comparable amount of time. The typical peak fluxes (which depend on the observing frequency) are of few $\\mu$Jy in the radio band with only a few OA reaching the mJy level. These values are consistent with the upper limits on the radio flux  of SN Ib/c observed at late times. We find that the OA radio number count distribution has  a typical slope $-1.7$ at high fluxes and a flatter ($-0.4$) slope at low fluxes with a break at a frequency--dependent flux. Our predictions of the OA rates are consistent with the (upper) limits of recent radio surveys and archive searches for radio transients. Future radio surveys like VAST/ASKAP at 1.4 GHz should detect $\\sim3\\times10^{-3}$ OA deg$^{-2}$ yr$^{-1}$, MeerKAT and EVLA  at 8.4 GHz should see $\\sim3\\times10^{-1}$ OA deg$^{-2}$ yr$^{-1}$. The SKA, reaching the $\\mu$Jy flux limit, could see up to $\\sim0.2-1.5$ OA deg$^{-2}$ yr$^{-1}$. These rates also depend on the duration of the OA above a certain flux limit and we discuss this effect with respect to the survey cadence. ",
        "introduction": "In the standard external shock model of Gamma Ray Bursts (GRBs) the afterglow emission is produced when the ultra relativistic jet is decelerated by the interstellar medium (Meszaros \\& Rees 1997). During this phase the bulk Lorentz factor $\\Gamma$  decreases with time while the beaming angle of the emitted radiation $\\Omega(t)\\propto 1/\\Gamma(t)^{2}$ increases. Moreover, there is a large consensus both for theoretical and observational reasons, that GRBs are jetted sources. The estimate of the jet opening angle can be derived from the time of the afterglow light curve steepening (the jet break time \\tjet) due to the Lorentz factor becoming $\\sim$1/\\th\\ (Rhoads et al. 1997).  Typical jet opening angles derived from the afterglow light curve breaks  (Frail et al. 2001; Ghirlanda et al. 2004) are clustered around \\th=0.05 radiants.  Therefore, a GRB which is observed off--axis, with a viewing angle \\thv$>$\\th, will be undetected as a prompt burst of $\\gamma$--ray photons because the prompt emission,  produced by material moving with a bulk Lorentz factor $\\Gamma_0\\sim 10^{2}-10^{3}$, is beamed within an angle 1/\\G$<$\\th. However, the  afterglow emission (in the optical/NIR and radio band) can be detected when the beaming angle of the radiation intercepts the line of sight, i.e.  1/$\\Gamma \\sim$\\thv. After this time the emission for an off-axis observer is the same that would be seen if \\thv$<$\\th.   Orphan Afterglows (OA) are GRBs seen off-axis detectable at any frequency, without the high energy $\\gamma$--ray counterpart. For this reason their study follows a different path than  normal GRBs (where it is the high energy trigger to initiate a follow up campaign to monitor the afterglow emission at different frequencies). OAs can be detected as transients through wide field deep surveys and they could be a considerable fraction of the population of detected transients (Rohads 1997, 2003;  Nakar, Piran \\& Granot 2002; Totani \\& Panaitescu 2002).  Considering the typical jet opening angles measured in GRBs (e.g. Frail et al. 2001; Ghirlanda et al. 2007) \\th$\\sim 3^{\\circ}$, for each GRB detected in the $\\gamma$--ray band  there should be $\\sim730$ (i.e. $\\sim2/$\\th$^2$) bursts pointing in any other direction. These are orphan afterglows.  Upper limits \\th$\\lsim$22$^\\circ$ (Levinson et al. 2002) or lower limits  \\th$\\gsim$0.8$^\\circ$ (Soderberg et al. 2006) on the typical opening angle of GRBs were inferred from OA searches. However, to date no orphan afterglow has been confirmed. Some transients have been identified as possible OA candidates through archival searches both at radio frequencies (e.g. see Murphy et al. 2013; Bell et al. 2011 for a summary of radio OA search results) and in the optical band (e.g. Rau, Greiner \\& Schwarz  2006; Malacrino et al. 2007; Zou, Wu \\& Dai  2007) but none of these has been confirmed as an orphan GRB afterglow. Even the very recent discovery (Cenko et al. 2013) of an optical transient by the Palomar Transient Survey (PTF) seems favor a ``dirty fireball\" or an untriggered bursts origin (i.e. both scenarios related to a GRB pointing towards the Earth) rather than to an orphan afterglow (Cenko et al. 2013).  A possible strategy (Soderberg et al. 2006; Bietenholz et al. 2013) for identifying off--axis GRBs is to observe a considerable radio emission from type Ib/c SNe at late times (i.e. years after the SN explosion). The estimate of the size of the SN event at late epochs allows us to verify if the SN had a relativistic jet. A handful of SNe  were found to have a bright radio emission: SN 2001em - (Granot \\& Ramirez Ruiz 2004), SN 2007gr (Paragi et al. 2010) and SN 2009bb (Soderberg et al. 2010), SN 2003gk (Bietenholz et al. 2013). However either subsequent revision of the radio observations (Soderberg et al. 2010a) or VLBI observations (Bietenholz, BArtel \\& Rupen 2010; Bietenholz et al. 2013) showed that these events are non--relativistic SN Ib/c. Only SN 2009bb could be considered as a transition event between the class of SNe and GRBs (Soderberg et al. 2010a).  If GRBs have a jet, OA should exist and they should be detected by wide field deep surveys. Here we consider the standard model of a uniform jet for GRBs but the detection rate of OA has been explored (Rossi, Perna \\& Daigne 2008) also for a universal structured jet model (Rossi, Lazzati \\& Rees 2002; Zhang \\& Meszaros 2002). \\begin{figure} \\includegraphics[width=8.5cm,trim=20 10 20 10,clip=true]{population.eps} \\caption{Viewing angle (\\thv) versus jet opening angle (\\th) of the simulated population of GRBs (G13). The solid line of equality separates GRBs pointing to the Earth (blue symbols - with \\th$\\ge$\\thv) from the bulk  of the population of GRBs not pointing to the Earth (black symbols - with \\th$\\le$\\thv) which can be detected as Orphan Afterglows. } \\label{fg1} \\end{figure}  If the  non detection of OA is due to the low sensitivity of past surveys, future deep and wide field transient surveys could succeed in detecting OA and,  as mentioned above, they could represent a considerable fraction of the detected transient population (e.g. Frail et al. 2012). However, one  challenge is how to disentangle the population of GRB OA from other possible sources producing similar transients. In the radio band, the detection of GRB afterglows (Chandra \\& Frail 2012) benefits from the fact that the brightest phase of the emission happens on timescales of few days (for GRBs observed on--axis), because initially suppressed by self--absorption. This allows us to detect and follow the GRB radio emission when the flux at higher frequencies is already below the  sensitivity of  available instruments. Late time observations of radio emission from GRBs is fundamental for constraining some key parameters of these sources (e.g. Berger, Kulkarni \\&  Frail 2003, 2004) and for estimating their global energetic (Frail, Kulkarni \\& Nicastro 1997; Frail, Metzger \\& Berger 2004; Shivvers \\& Berger 2011; Sironi \\& Giannios 2013).  We are now entering a new era of radio surveys. In preparation for the Square Kilometre Array (SKA, Carilli \\& Rawlings, 2004), that will represent a giant leap forward in survey depth at all GHz and sub-GHz frequencies, there are several pathfinders that will become operative within the next few years. In the sub-GHz regime, both the Low Frequency Array for Radio-astronomy (LOFAR, van Haarlem et al. 2013) and the Murchinson Widefield Array (MWA, Bowman et al. 2013, Tingay et al. 2013) will scan the sky with unprecedented survey speed, thanks to their field of view. Nonetheless as far as the search for OA is concerned, these will be likely limited to the exceptionally bright sources. In the GHz band, the most promising telescopes -in terms of dynamic range and fields of view- for the systematic search for transients are the Australian SKA-Pathfinder (ASKAP, Johnston et al. 2007), the Aperture Tiles In Focus (AperTIF) experiment, a wide field upgrade to the Westerbrok Synthesis Radio Telescope (Verheijen et al. 2008), and the MeerKAT telescope. While the MHz telescopes are already operative, the GHz SKA pathfinders will take first light in the next 2--5 years.  In order to study the detectability of OA with on--going and future surveys it is necessary to know the emission properties (timescales and flux level) of the population of OA which depend on the properties (energetics, distance scale, jet opening angle) of the population of GRBs. In this paper we use a population synthesis code PSYCHE (Ghirlanda et al. 2013a; 2013b)  summarised in \\S2 which has already been used to predict the detectability of on--axis GRB afterglows in the radio band (Ghirlanda et al. 2013b). Here we explore for the first time with such a code the radio emission properties of Orphan Afterglows (\\S3) and compare with the current limits on the detection rates at radio frequencies (\\S4). We explore the detectability of OA by future radio surveys in \\S5. ",
        "conclusions": "Orphan afterglows are GRBs whose emission is detectable only during the afterglow phase (at optical/NIR and radio frequencies). Their prompt $\\gamma$--ray emission is unobservable because the viewing angle \\thv\\ is larger than the jet opening angle \\th\\ (off--axis GRBs). In these events the afterglow emission becomes observable when the bulk Lorentz factor, which is decreasing during the afterglow phase, becomes $\\Gamma \\sim 1/$\\thv. After this time, which represents the peak of the OA light curve, the emission is similar to that for an observer within the jet opening angle.  OA make up a majority of the population of GRBs. However, none have been observed so far, do to their lack of a prompt emission trigger. Their detection is possible as transients in deep/wide field surveys. However, so far no detection of OA has been confirmed by searches in archival optical/radio observations. In current and future surveys OA might  represent a considerable fraction of detected transients.  In this paper we have used the results of a population synthesis code for GRBs (G13) that simulates the entire population of GRBs including off-axis events and is anchored to reproduce some observational constraints of the population of  GRBs detected \\fe\\ and \\cgro\\ with particular emphasis on the constraints given by the BAT6 complete \\sw\\ sample (Salvaterra et al. 2012).  We have explored the properties of the population of off--axis GRBs (see Fig.\\ref{fg1} - black symbols) in terms of their radio emission. We have computed the radio flux density of the OA population (representing $\\sim$97\\% of the entire GRB simulated population) at the time when the OA light curve reaches its peak (Fig.\\ref{fg2}) which is of the order of few years after the prompt trigger. However, the lack of any prompt emission (i.e. $\\gamma$--ray trigger) in OA, requires to compute the timescale of their duration which, given the typical rise/decay long--term evolution of the afterglow flux, can be of the same order of the peak time. This suggest that OA in the radio band should be slow transients.  We have constructed the cumulative flux distribution at different radio frequencies (GHz and MHz) that shows a high flux tail with a slope consistent with $-1.7$ (Fig.\\ref{fg3}) and a break at a frequency dependent flux below which the slope becomes flatter ($-0.4$). This is due to the combination of the \\th\\ distribution of the population of simulated bursts with the viewing angle \\thv\\ probability function: at approximately 10$^{\\circ}$ the product of the \\th\\ log--normal distribution (resulting from the population code of G13) and of the probability density of \\thv\\ is maximised. This accounts for the slope change of the flux distributions in Fig.\\ref{fg3}, the break of the flux distribution is shifted to lower fluxes at lower frequencies because the radio emission is in the self absorbed regime of the synchrotron spectrum.  In general, from Fig.\\ref{fg3} we note that very bright OA with flux at the 1 mJy level are rare. This is the central flux of the current radio searches/surveys that have been searching for OAs. Our population is consistent with the limits given by these surveys. Totani \\& Panaitescu (2002) derived the flux distribution of orphan afterglows based on 10 bright GRB afterglows. Our estimates are consistent in the bright flux end with their but we predict a lower number of OA at low fluxes (below the break of our flux distribution). This is due to the modelling, in our case, of the GRB jet opening angle distribution (G13).  In G13 we derived that $\\sim$0.3\\% of SN Ib/c may harbour a GRB, i.e. in other words the great majority of the putative progenitors of GRBs do not produce a relativistic jet. When compared to the combined samples of S06 and B13, which comprise 112 SN Ib/c, there were indeed only two detections and we statistically expect none of them to be an off-axis GRB. Both SN 2001em and SN 2003gk were successively showed not to expand at relativistic velocity, arguing against their being genuine OAs. Finally, we showed (in Fig.~\\ref{fg5}) that all upper limits on SN Ib/c radio emission at late times are indeed consistent with the distribution of GRB OAs in the plane $F_{\\rm peak}$ -- \\tpeak."
    },
    "1402/1402.3940_arXiv.txt": {
        "abstract": "Temporal and spectral variations of black hole candidates during outbursts have been reported in several publications. It is well known that during an outburst, source becomes softer in first few days, and then returns to hard states after a few weeks or months. In the present paper, we show variation of Comptonization Efficiency (CE), defined to be the ratio of number of power-law photons and number of soft photons injected to into the Compton cloud, as a function of time in several outbursts. Since power-law photons are generated through inverse-Comptonization of the intercepted soft photons, CE is a measure of geometry of Compton cloud. Our investigation indicates that all outbursts start with a large CE and becomes very small after a few days, when the Compton cloud becomes small enough to intercept any significant number of soft photons. CE returns back to a larger value at the end of outburst. We show co-variation of count rates, frequency of quasi-periodic oscillations (QPOs), photon index and CE and establish a general trend of disk geometry variation during outbursts in all systems under consideration. ",
        "introduction": "It is well known that black hole accretion flows in a compact binary system typically consist of a Keplerian disk which emits soft or low energy photons and a hot Compton cloud which inverse Comptonizes soft photons into high energy photons \\citep{sun85}. When a black hole spectral state changes from hard to soft and vice versa, Compton cloud must change its shape and temperature \\citep{C95,t95}: In a soft state, when the accretion rate in the Keplerian disk is high, Compton cloud is smaller and cooler, while in a hard state, when Keplerian disk rate is very low, Compton cloud is larger and hotter. There are various models of the Compton cloud in the literature, ranging from a hot Corona, magnetic corona, to post-shock region of an accreting low angular momentum flow which surrounds the standard Kelerian disk \\citep{H91, w91, e97, J00, MF01, z03, N11, O05, li05, C95, CM00, rao00, M06}. Along with spectral variabilities, a black hole candidate shows temporal variations also, which is sometimes quasi-periodic \\citep{M99, C08, C09}. There are various explanations of quasi-periodic oscillations in literature ranging from disc-seismology to oscillation of post-shock region \\citep{k05, ax05, CM00}.  Outburst sources are ideal candidates to study changes in the size of a Compton cloud. This is because object is known to be in a hard state at the beginning of an outburst but in a matter of few days to a few weeks, it changes its states to several other states, thereby giving us a unique opportunity to study size variation of Compton cloud very easily. From the spectral analysis, it appears that following sequence is typically followed by most, if not all, outburst sources: hard $\\rightarrow$ hard-intermediate $\\rightarrow$ soft-intermediate $\\rightarrow$ soft $\\rightarrow$ soft-intermediate $\\rightarrow$ hard-intermediate $\\rightarrow$ hard) \\citep{C08, C09, B10, n12}.  Recently, \\citet{P11} showed that average size of Compton cloud in the variable source GRS 1915+105 changes in a very well defined way as it transits from one variability class to another. They computed a quantity called Comptonizing Efficiency (CE) which is the ratio of number of power-law photons to the number of blackbody photons in the spectrum at a given instant. They show that CE is very small in softer classes and larger for harder classes. Within some of classes, there are evidences of rise and fall of count rates and spectral slopes in a matter of a few seconds. Since number of hard photons depends on optical depth of Compton cloud, we clearly see a change in optical depth of the Compton cloud in that time scale. Furthermore, we see that one variability class goes to another adjacent class whose CE is nearest.  In the present paper, we study variation of Compton cloud size during outbursts of several black hole candidates. We show that outbursts typically start with a large Comptonization efficiency, i.e., a large sized Compton cloud with a poor soft photon source. In rising phase, this cloud becomes progressively smaller and smaller on a daily basis, till it became minimum when the object went to a soft state. Of course, soft photon source gets stronger (soft photon intensity rises) and as a result, the cloud size shrinks (QPO frequency, assumed to be the oscillating frequency of the cloud, rises). In declining phase of an outburst, the trend is exactly reversed. In this paper, we deal with black hole candidates which exhibited outbursts, such as, H1743-322, GX 339-4, 4U 1543-47, XTE J1118+480, XTE J1859+226, GRO J1655-40 and XTE J1550-564. In the next Section, we briefly discuss the objects of our study. In \\S 3, we present the procedure of our analysis. Here, we compute Comptonization Efficiency (CE) of all these outbursts both in rising and declining phases. We show that CE changes with time very rapidly. In general, softer states were found to have a very low CE and the harder states (at the beginning of rising phase and at the end of declining phase) were found to have a very high CE. In \\S 4, we present results and interpret them using physical models. Finally, in \\S 5, we draw our conclusions. ",
        "conclusions": "In this paper, we analyzed several black hole candidates which exhibit outbursts and shown the time dependence of Comptonizing efficiency, QPOs and the spectral index. We tried to understand the results using the TCAF model of \\citep{C95} and its time varying form, namely, propagatory oscillating shock (POS) model of the outbursts \\citep{C08, C09, D10, DD10}, though any model which relies on soft photons and inverse Comptonization would be fine, since CE depends on the ratio of Comptonized photons and soft seed photons. In TCAF model, a standard Keplerian disk is surrounded by a faster moving low-angular momentum flow (sub-Keplerian) which produces a centrifugal pressure supported shock where the flow is puffed up and produces the so-called Compton cloud to inverse Comptonize intercepted soft photons coming from the Keplerian disk. Oscillation of post-shock region or CENBOL causes low frequency QPOs in black hole candidates.  In the backdrop of this model, our goal is to study how the optical depth of the CENBOL changes with time in a generic outburst source. If we had computed hardness ratio or HR, where the soft and hard photons are counted using the same fixed energy bins, a comparison of its behavior from one object to another would not be possible. This is because seed photons of one black hole could be Comptonized photon for another. On the other hand, CE as defined by us characterizes soft and hard photons objectively and as such does not depend on mass of the black hole or its accretion rate. Thus a comparison is possible. It is true that we restrict ourselves from 0.1 to 40 keV photons. However, for stellar mass black holes, these boundary values are far away from relevant energies in which respective photons are important.  We came to a conclusion that generally speaking, all outbursts start and end with a large Compton cloud size, though not necessarily with the highest optical depth. As the outburst progresses, CE becomes minimum at peak of an outburst i.e., size of CENBOL becomes very small. We find that CE in outburst sources may vary from almost $\\sim 0.0$ to $\\sim 3$\\%. In contrast, variable source GRS 1915+105 has CE between $0.005$ to $0.8$ and has relatively high luminosity (even after factoring out effects of mass of the black hole), suggesting that it is in a soft-intermediate state which usually appears after the peak of a possible outburst. If so, in future, this source may slow down its activities when the viscosity in this system is reduced as in any other outbursting source.  We would like to stress that CE computed by us counts soft photons from $0.1$ keV to an upper energy limit automatically detected by our fitting method. If we concentrated on photons in the observed $2-40$ keV range and counted the number of soft photons, we would have obtained a different, in fact, much higher CE. This is because we would have grossly underestimated number of soft photons. In Fig.~\\ref{figfrac}, we show a comparison of CE computed by our method with `fracsctr' obtained from {\\it  simpl} model \\citep{st11} for GX 339-4 data. {\\it simpl} is an empirical model of Comptonization which computes fraction of photons from an input seed spectrum which are scattered into a power-law component. Fracsctr is fraction (in the scale up to 1.0) of input seed photons is Comptonized. Black and gray points are for the 2007 and 2010 outburst data respectively. Black stars and gray circles represent variations of fracsctr with CE when soft photons are taken from 0.1-40 keV CE(\\%) range. One generally sees that one increases with the other. From Fig.~\\ref{figfrac}, we can see that fracsctr value varies around 0 to 0.8 this means 0-80 $\\%$ photons are interacting with hot electron cloud. But if we compare with CE then only 0.005-0.35 $\\%$ of photons are interacting with hot electron cloud which is more realistic. Computed CE values do support relative shape of CENBOL which is expected to intercept less than a few percent of soft photons from the Keplerian disk \\citep{C95}.  \\begin{figure} \\begin{center} \\includegraphics[width=7.0truecm,angle=270]{fracsctr_new.ps} \\caption{Variation of fracsctr with CE(\\%) for 2007 and 2010 outbursts of GX 339-4 data. Black stars and gray circle points respectively represent data of 2007 and 2010 outbursts. This plot shows variation of fracsctr with CE(\\%) computed in 0.1-40 keV range.} \\label{figfrac} \\end{center} \\end{figure}  Another important issue is to correlate spectral and timing data. We mentioned that CE is higher for harder states. This is because location and height of shock wave are large. We have also mentioned that low frequency QPOs are due to oscillations of these shock waves and frequency is inverse of the infall time scale from post-shock region \\citep{CM00} till the inner sonic point. As such, larger shock radius means smaller QPO frequency, which is exactly what we observe! If we combine these results, then TCAF model would predict that CE would be larger when the QPO frequency is smaller and vice versa. In Fig.~\\ref{qce}, we show measured QPO frequencies as a function of our CE for all the outbursts discussed in this paper. We find that CE is inversely correlated with QPO frequencies, although the relationship is not very tight in some cases. This could be because of presence of radio jets observed in all these outburst sources \\citep{BR02,Cas10,Cor10,K05,B04,Cor08, Mig07,H95,K03}. While CE is influenced by electron clouds at base of a jet (post-shock region in TCAF solution), QPO frequency is not directly affected by the jet, but only by location of oscillating shock. This could be a primary reason why some outbursters do not show a tight correlation. However, this aspect requires further investigation.  \\begin{figure} \\begin{center} \\includegraphics[width=7.0truecm,angle=0]{qce_9_new.ps} \\caption{Variation of QPO frequency with CE(\\%) for all outbursts. The general trend is that the frequency increases as CE is decreased.} \\label{qce} \\end{center} \\end{figure}  It is also interesting to study variation of CE with luminosity in analogy with hardness ratio-intensity plots in the literature \\citep{f04}. Since hardness ratio is obtained from a fixed range of X-ray photons, independent of mass of black hole, it is difficult to interpret what a conventional colour-colour diagram really means. In Figs.~\\ref{qdia}a, and \\ref{qdia}b, we plot CE vs. luminosity for both outbursts of XTE 1550-564 and GX 339-4 respectively. In Fig.~\\ref{qdia}(a), dark filled boxes represent variation of CE(\\%) with natural logarithm of luminosity in Eddington unit during 1998 outburst and gray filled circles represent same for 2000 outburst. In Fig.~\\ref{qdia}(b), black filled boxes represent 2007 outburst of GX 339-4, while gray filled circles represent the same for 2010 outburst of GX 339-4. For all plots, arrows of corresponding colour are provided to understand evolution of data during outbursts. Figures provide an idea about the relation between geometric size of the electron cloud and total luminosity during outburst. In case of XTE 1550-564 (Fig.~\\ref{qdia}a), we observe that both outbursts started and ended with high CEs, though final value of CE is lower than its initial value. However, this could be due to incompleteness of observation during onset of the outburst. Outburst of 2000 is clearly weaker and short lived. In both the cases, there is a hysteresis effect in that paths in rising and declining phases are not the same. This is expected, since formation time and disappearance time of a Keplerian disk by viscous effects are not identical \\citep{G13}. In case of GX 339-4, both Figures are generally overlapping and thus both were of roughly equal strength. We note that CE is not necessarily highest at beginning or end in all these cases, though it is close to highest values achieved during outburst. This may be because CE is sensitive to the optical depth and not just physical size of the CENBOL. This is also the reason why QPO frequency variation (sensitive to the size of CENBOL) is not tightly correlated with CE. An important result is that CE of GX 339-4 is much lower as compared to that of XTE J1550-564. Since with spin the size of the CENBOL shrinks at least by a factor of two \\citep{C96}, the degree of interception would also be reduced by the same factor. It is well known that spin of XTE J1550-564 is moderate \\citep{st11} while the spin of GX 339-4 is very high \\citep{kol10}. It is possible that we are seeing effects of large spin in GX 339-4 in our calculation also.  \\begin{figure} \\includegraphics[width=6.35truecm,angle=270]{qdia_1550_new.ps} \\includegraphics[width=6.35truecm,angle=270]{qdia_gx_new.ps} \\caption{Variation of CE(\\%) with $(L/L_{Edd})$ for (a) 1998(black) and 2000(gray) outburst of XTE J1550-564, (b) 2007(black) and 2010(gray) outburst of GX 339-4.} \\label{qdia} \\end{figure}   In order to show that soft photons beyond $dbb_e$ really do not contribute significantly, we recomputed the above values with soft photons till $40$ keV. We do not find any significant difference in the result. So we believe that our results with dynamically obtained black body photon energy range yields sufficiently accurate results.  From the Table:~3, we see that range of CE is different for different objects. Duration of outbursts, QPO frequency range etc. are also found to be different. There is also a considerable scatter in CE which may be due to outflows from the CENBOL. All these require a more thorough analysis for a unifying understanding of the outbursts. This will be addressed in our future publications. Another way to improve the result is to use TCAF model itself to fit the spectra and obtain blackbody energy range more accurately. A third, and more serious point is to consider spin of black hole which will reduce CE by virtue of smaller CENBOL size for higher spins. These aspects will be looked into in future."
    },
    "1402/1402.4567_arXiv.txt": {
        "abstract": "We have developed a new formulation to obtain self-gravitating, axisymmetric configurations in permanent rotation. The formulation is based on the Lagrangian variational principle, and treats not only barotropic but also baroclinic equations of state, for which angular momentum distributions are not necessarily cylindrical. We adopt a Monte Carlo technique, which is analogous to those employed in other fields, e.g.  nuclear physics, in minimizing the energy functional, which is evaluated on a triangulated mesh. This letter is a proof of principle and detailed comparisons with existing results will be reported in the sequel, but  some test calculations are presented, in which we have achieved an error of $O(10^{-4})$ in the Virial relation. We have in mind the application of this method to two-dimensional calculations of the evolutions of rotating stars, for which the Lagrangian formulation is best suited. ",
        "introduction": "After the invention of the Henyey method~\\citep{henyey64}, our understanding of the structures and evolutions of spherically symmetric stars in the absence of rotation (and/or magnetic fields) have been advanced exponentially over the years by many researchers\\citep{maeder00}. Although the stellar evolution is a time-dependent problem, it is actually so slow in most stages compared with the dynamical time scale that it is well approximated by a collection of time-independent configurations in equilibrium with gradually changing parameters such as the central density and temperature as well as chemical compositions.  In the presence of rotation, however, the same method cannot be employed. One of the difficulties in calculating rotational equilibria is the fact that we do not know a priori what the distribution of angular momentum looks like in the stellar interior. This is indeed the case even for barotropes, in which matter pressure is a function of density alone, although it is mathematically shown in this case that the rotational equilibria have cylindrical distributions of specific angular momentum. When the angular momentum takes constant values on concentric cylinders, there exists a first integral and there are a variety of numerical treatments~\\citep{ostriker, hachisu86}. The realistic equation of state, however, is not barotropic but baroclinic, in which pressure depends on density, entropy and chemical compositions. It is hence a non-trivial issue how to make an appropriate sequence out of an ensemble of rotational configurations even if they are somehow obtained~\\citep{uryu94, espinoza13}.  As a matter of fact, in the application of rotational equilibria, we would then need to know the temporal evolutions of thermodynamics quantities as well as angular momentum distribution, which would not be easy in multi-dimensions, since fluid elements composing rotating stars could change their positions non-radially in complicated ways even without convection. Such gradual evolutions that occur on the secular time scale are best treated by the Lagrangian formulation just as in the spherically symmetric case. This is the idea underlying our method.  It is observationally well known that massive stars rotate at a surface velocity of $\\sim 200\\ {\\rm km\\ s}^{-1}$ on the main sequence~\\citep{fukuda82}. Although the angular momentum in  stellar interior is largely unknown, it will not be a small value. In fact, centrifugal force by such rotations can play a crucial role in the advanced evolutionary stages of massive stars~\\citep{endal76}. In particular, various (magneto-) hydrodynamical instabilities induced by rotation have been intensively and extensively investigated over the years by many researchers (see e.g. \\citealt{maeder00, maeder13}). Full-fledged instabilities may produce turbulence even in the radiative zones~\\citep{mathis13}. Since such turbulences are expected to be anisotropic in general, the time-averaged states will be approximated by some baroclinic rotational equilibria.  In this letter, we propose a novel idea to obtain rotational equilibria numerically for both barotropes and baroclines\\footnote{Non-barotropic equations of state are collectively referred to as baroclinic equations of state.}, which is well suited for the application to the stellar evolution calculations. It is based on the Lagrangian variational principle: an energy functional is defined in such a way that its stationary point should correspond to a rotational equilibrium. It is then approximated on a triangulated mesh and minimized by moving grid points, a Monte-Carlo technique quite analogous to what is commonly used in other fields, e.g. in nuclear physics to obtain deformed nuclei. We demonstrate its nice performance in some test models although this letter is meant to be a proof of principle and detailed validations will be presented in the sequel. ",
        "conclusions": "It is emphasized that the establishment of the formulation that can handle not only barotropic but also baroclinic rotational equilibria is itself a major achievement, since such attempts are not many in the literature~\\citep{uryu94, espinoza13}.  Note also that all previous 2D formulations are Eulerian, whereas our formulation is Lagrangian in 2D and we can trace the potentially complex movement of each fluid element during the quasi-static stellar evolutions. Its advantage in investigating the evolutionary series of rotating stars should be obvious. Nuclear reactions occur locally in each fluid element and can be treated by network calculations on each node in our method. Radiative energy transport will be handled by solving a diffusion equation on the triangulated mesh. Angular momentum transfer and mixing of elements by various (magneto) hydrodynamical instabilities~\\citep{maeder13} may be also approximated as diffusions on the Lagrangian coordinates just as in the current 1D stellar evolution models~\\citep{ maeder00}.  The meridional circulation, which is neglected in the current formulation, is another important issue\\citep{LNP}. Since it is rather slow~\\citep{espinoza13}, it is insignificant in force balance \\citep{LNP}. It plays a major role, however, in the re-distribution of angular momentum and chemical abundance and, as a consequence, drives the evolution of rotating stars. Our strategy to treat this is the following: working on the time scale that is much longer than the dynamical time scale, we will incorporate this re-distribution of angular momentum and elements by by solving the same transport equations on the triangulated mesh as employed in the 1D models~\\citep{LNP}.  The biggest challenge is the treatment of convectively unstable configurations. Whereas the energy functional is {\\it minimized} in our formulation, convectively unstable equilibria do not correspond to the minimum but to a stationary point (or a saddle point) of the energy functional. A naive application of the formulation to such configurations will lead to a continued distortion of the mesh in the convection zone. A couple of ideas are currently being tested and will be reported elsewhere.  As for the numerical techniques, the Monte Carlo method employed in this letter is rather slow and it takes days to get the minimization done at present. Although this may appear to be a concern in the application to the stellar evolution, it should be emphasized that the minimization is achieved in much shorter time if the initial configuration is very close to equilibrium, which is actually the case in the stellar evolution calculations.  Appropriate triangulations, which are not touched in this letter, will become important in the application of the method to stars in advanced evolutionary stages, in which the stars develop a core and mantles with large density contrasts. It may be necessary to extend the formulation so that it could handle multi-layers. Possible applications are not limited to these stars but will be extended to e.g. compact stars, proto-stars and planets.  Incorporation of magnetic fields and/or general relativity should be considered in due course. Our new formulation will be a real break-through then."
    },
    "1402/1402.5082_arXiv.txt": {
        "abstract": "The relation between the stellar mass and size of a galaxy's structural subcomponents, such as discs and spheroids, is a powerful way to understand the processes involved in their formation. Using very large catalogues of photometric bulge+disc structural decompositions and stellar masses from the Sloan Digital Sky Survey Data Release Seven, we carefully define two large subsamples of spheroids in a quantitative manner such that both samples share similar characteristics with one important exception: the `bulges' are embedded in a disc and the `pure spheroids' are galaxies with a single structural component. Our bulge and pure spheroid subsample sizes are 76,012 and 171,243 respectively. Above a stellar mass of \\Critmass{}, the mass-size relations of both subsamples are parallel to one another and are close to lines of constant surface mass density. However, the relations are offset by a factor of 1.4, which may be explained by the dominance of dissipation in their formation processes. Whereas the size-mass relation of bulges in discs is consistent with gas-rich mergers, pure spheroids appear to have been formed via a combination of `dry' and `wet' mergers. ",
        "introduction": "The structure of galaxies emerges from a complex process of hierarchical mass assembly in which mergers and other processes such as dissipation play key roles. Discs and spheroids contain the bulk of the light (and thus stellar mass) of galaxies, and their scaling relations encode valuable information about their formation and evolution \\citep[e.g.][]{fall80, djorgovski87, burstein97, mo98}. Discs and spheroids have been most often distinguished on the basis of photometric data from which light profiles are measured and morphological decompositions are performed \\citep[e.g.][]{freeman70,boroson81,kent85,dejong96,peng02,simard02,simard11,lackner12,meert13}. Under this observational classification paradigm, spheroids are parametrically described by S\\'ersic/de Vaucouleurs light profiles, whereas discs' profiles follow an exponential decline. However, this relatively simple picture is changing with the greater availability of detailed internal kinematics data, e. g., ATLAS$^{\\rm 3D}$\\citep{krajnovic12}. Through comparisons between photometric and kinematical data, fast rotators have been associated with disc galaxies and slow rotators with dispersion-supported spheroids. The fast/slow rotator classification is also closer to the definition of structure in numerical simulations \\citep[e.g.][]{abadi03,governato07,hopkins09a,agertz11,brooks11,scannapieco11}. High angular momentum particles in these simulations are associated with a disc, and spheroids comprise the low angular momentum particles.  An extensive body of theoretical and observational studies \\citep[e.g.,][]{bender92,ciotti01,shen03,robertson06,ciotti07,scarlata07,hopkins09a,nipoti09,auger10,nipoti12,shankar13,taranu13} has shown that merger histories and hydrodynamical processes such as dissipation shape the properties of spheroids. Much of this work has focused on reproducing the mass-size relation. This fundamental relation describes how the sizes of galaxies scale with their stellar mass. The relation is a projection of the Fundamental Plane \\citep{djorgovski87}, an inherent relation for all dispersion supported galaxies and is typically very tight \\citep[e.g.][]{lauer07,kormendy09}.  However, there are interesting trends that remain to be studied in greater detail. For example, \\citet{shen03} found that the merger of two similar late-type galaxies produced an apparent flattening of the mass-size relation of the remnant spheroid and that disc instabilities were responsible for growing the bulge. \\citet{shankar13} found a similar flattening, but they attributed it to the lack of gas in the merging galaxies. There is also evidence that bulges are not present-day elliptical galaxies embedded in discs \\citep{graham08,gadotti09,gadotti09-2,laurikainen10} as they exhibit different mass-size relations. %  The gas content of a merger progenitor can determine the size of the virialized remnant as it controls the amount of energy that can be lost as a result of dissipation. Gas-rich (`wet') mergers tend to dissipate orbital energy, which shrinks the effective size of the system as the gas cools. The light profile of such a galaxy will have a higher stellar surface density due to the increasing central concentration of stars \\citep[][and references therein]{hopkins09b}. However gas-poor (`dry') mergers maintain this energy in stellar orbits, keeping the galaxy puffed up. Violent relaxation can mimic dry mergers by ejecting stars into the outer regions of galaxies, making the light profile extend to larger radii. Discs can regrow and be replenished around these merger remnants both by diffuse infall from the surroundings and by gas accreted from satellite galaxies \\citep[e.g.][]{guo11}. Although one might expect discs to be destroyed in the merging process itself, both \\citet{robertson06} and \\citet{hopkins09b} (amongst others) demonstrated that discs can quickly reform during the merger with sufficiently gas-rich progenitor galaxies. The bulge is then built very slowly with little angular momentum, keeping it small. Finally, the bulge component of disc galaxies can also form via secular instabilities in the disc, transferring some disc material into the bulge component \\citep[][and references therein]{courteau96, shen03}.  \\citet{shankar13} have recently demonstrated with semi-analytic models of galaxy evolution \\citep{guo11} that the mass-size relation flattens for spheroidal galaxies formed by \\emph{dissipationless} processes below a critical mass of \\Critmassel{}. This result differs from their sample of observations of early-type galaxies from the Sloan Digital Sky Survey (SDSS). The pure power law description of their data was suggested to be the result of \\emph{dissipation} of energy due to wet mergers. The flattening effect seen below a critical mass is believed to be an event of dry mergers as the excess energy puffs up the spheroidal component of the galaxy.  \\citet{hopkins09b} also simulated several mergers of different masses and different gas contents. They found that when dry mergers are turned off (i.e. only wet mergers); the sizes are much smaller than expected, although the mass-size relation still remains more or less a power law. When gas was treated like stars (i.e. no dissipation), a power law similar to the `no dry merger' case resulted; however, these sizes are larger than those from wet mergers. These two studies show that the size of the bulges in gas-rich galaxies seems to be smaller than their dry-merging counterparts. However, both pieces of work disagree on how strong the effects of dissipation are as their mass-size relations are inconsistent with each other, especially with the flattening seen by \\citet{shankar13} at low masses.  In order to better understand the effects of dissipation on the formation of the spheroidal components of galaxies, we study the mass-size relations of two very large spheroidal subsamples, namely bulges and pure spheroids, and compare them with expectations from theoretical models. The careful, quantitative selection of our subsamples, derivation of spheroid masses and the characterization of the mass-size relation are described in Section~\\ref{data}.  Our results are discussed in Section~\\ref{results}. The cosmology adopted throughout this paper is ($H_0, \\Omega_{m}, \\Omega_{\\Lambda}$) = (70 km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7). ",
        "conclusions": "Using large catalogues of bulge+disc structural parameters and stellar masses for galaxies in the SDSS DR7 Legacy Area, we quantitatively defined two subsamples of galaxy spheroids to distinctly select pure spheroids and bulges with identical properties other than bulges being embedded in disks. These two subsamples are used to study the relation between their size and stellar mass. We found that these subsamples had distinct mass-size relations, which indicates that bulges are not simply present day, pure spheroids wrapped in discs. At stellar masses above \\Critmass{}, the relations are parallel but shifted by a factor of 1.4 in size at a fixed stellar mass. Bulges follow a mass-size relation consistent with a `wet' merger model, and the relation for pure spheroids indicates a combination of `wet' and `dry' mergers. At stellar masses below \\Critmass{}, both relations exhibit a flattening in size with decreasing mass that is not reproduced at all by theoretical models. A comparison between our observed relations and theoretical models suggest that dissipation leaves a distinctive signature on the mass-size relations of bulges and pure spheroids."
    },
    "1402/1402.6518_arXiv.txt": {
        "abstract": "{We present thirteen new transit light curves for the WASP-1~b exoplanet. Observations were acquired with 0.5 - 1.2-m telescopes between 2007 and 2013. Our homogeneous analysis, which also includes the literature data, results in determining precise system parameters. New values are in agreement with those reported in previous studies. Transit times follow a linear ephemeris with no sign of any transit time variations. This finding is in line with the paradigm that Jupiter-like planets on tight orbits are devoid of close planetary companions.}{planetary systems -- stars: individual: WASP-1 -- planets and satellites: individual: WASP-1~b} ",
        "introduction": "Observations of transiting exoplanets offer a wealth of knowledge about extrasolar planetary systems. Photometric and spectroscopic datasets, when combined together, allow the planetary radius and exact planetary mass to be determined. These quantities give the mean planetary density, which is a clue to the internal structure of a planet (\\eg Nettelmann \\etal 2010). Observations of occultations of transiting planets (\\eg Snellen \\etal 2009) constrain an orbital eccentricity, which is usually poorly determined from radial velocity (RV) measurements alone. These observations also allow one to determine chemical properties of the planetary atmosphere with the spectrophotometry technique (\\eg Charbonneau \\etal 2008). Observations of the Rossiter--McLaughlin effect provide the sky-projected angle between the stellar spin axis and the planetary orbital axis (\\eg Queloz \\etal 2000). This information gives important insight into theories of planet formation and dynamical evolution.  The architecture of planetary systems with a transiting planet may be studied with the RV and transit timing methods to search for additional planets. The former technique has already brought discoveries of companions on wide orbits (\\eg Bakos et al. 2009). The latter method, based on the detection of transit time variations (TTVs,  \\eg Agol \\etal 2005, Holman \\& Murray 2005), is sensitive to terrestrial-mass planetary companions close to low-order mean-motion resonances (MMRs, Steffen \\etal 2007). In such configurations, the mutual gravitational interaction of planets may generate deviations from strictly periodic transit times. The terrestrial-mass planets are usually below the current RV sensitivity limit which makes the TTV method, {\\bf in addition to gravitational microlensing}, the unique discovery tool. So far, the TTV technique has been successfully applied to data from the space-based {\\it Kepler} transit survey (Borucki \\etal 2010). In transiting multi-planetary systems, analyses of TTVs have allowed one to confirm or validate the planetary nature of many candidates (\\eg Holman \\etal 2010). In single transiting planet systems, dynamical analyses of observed transit times have led to discoveries of unseen planetary companions (\\eg Ballard \\etal 2011), some of them confirmed independently by RV measurements (Barros \\etal 2014).  Statistical studies of hot-Jupiter candidates (planets with masses of the order of 1 Jupiter mass, $M_{\\rm{Jup}}$, and orbital periods shorter than 5 days) in the {\\it Kepler} sample show that these planets accompanied by nearby low-mass planets are rare (Steffen \\etal 2012, Steffen \\& Farr 2013). In contrast, planets on slightly wider orbits, so called warm Jupiters, Neptunes, or super Earths, are observed to form multi-planetary systems with planetary pairs close to MMRs (\\eg Ming \\etal 2013). The relative isolation of hot Jupiters could be the result of the initial formation conditions and dynamical evolution of their planetary systems (Kley \\& Nelson 2012). Theoretical studies show that the system architecture may depend on many factors, including the mass ratio of planets relative to the central star, the mass of the protoplanetary disk, and the mechanism of the planetary migration (Goldreich \\& Schlichting 2013). Thus, the detection or a lack of companion planets close to hot Jupiters may put forward observational constraints on the theory of planet formation (\\eg Ogihara \\etal 2013; see also Steffen \\etal 2012 for a detailed discussion).  It has been shown that a terrestrial-mass planet perturbing a hot-Jupiter gas giant in an orbital resonance can produce a TTV signal of an amplitude of a few minutes (Steffen \\etal 2007). Such a signal can be detected using timing observations acquired with 0.5--1.0-m telescopes. Being encouraged by this finding, we have organized observing campaigns for selected transiting planets to search for their possible TTVs (Maciejewski \\etal 2011a). So far, we have put constraints on the upper mass of additional planets in WASP-3 and TrES-3 systems (Maciejewski \\etal 2013b, Va\\v{n}ko \\etal 2013), refined transit ephemerides for some transiting planets (Maciejewski \\etal 2011c, 2013c, Raetz \\etal 2014), studied physical properties of the WASP-10~b planet and the activity of its host star (Maciejewski \\etal 2011b), and finally found a hint of orbital perturbations for the WASP-12~b planet (Maciejewski \\etal 2013a). In this paper, we study the transiting exoplanet WASP-1~b (Collier Cameron \\etal 2007). Our observing campaign, spanning six years, resulted in refining system parameters and putting upper-mass constraints on the second planet in the system.  The WASP-1 star (GSC 02265-00107, $V=11.31$ mag) with its transiting planet belongs to the first planetary systems discovered by the SuperWASP survey (Pollacco \\etal 2006). The planetary nature of its companion was confirmed with RV measurements acquired with the SOPHIE spectrograph (Bouchy \\etal 2006). The planet was found to be bloated with a sub-Jovian mass and the radius 1.4 times greater than that of Jupiter. It orbits a dwarf star of the F7 spectral type on a 2.5-d orbit with the semi-major axis of 0.038~AU. The system parameters were refined by follow-up observations by Charbonneau \\etal (2007) and Shporer \\etal (2008) and a re-analysis of Southworth (2008; 2010). Szab\\'o \\etal (2010) refined the transit ephemeris with new transit observations acquired with a small telescope.  The host star was the subject of a detailed spectroscopic study by Stempels \\etal (2007). Precise RV measurements, obtained throughout a transit by Simpson \\etal (2011) and Albrecht \\etal (2011), show that WASP-1~b belongs to the group of misaligned planetary systems with the almost polar orbit. Wheatley \\etal (2010) observed an occultation of WASP-1~b with the Spitzer Space Telescope. Their results show that there is a strong temperature inversion in the planet's atmosphere. Furthermore, the phase of mid-occultation was found to be consistent with a value of 0.5. This finding suggests that the planetary orbit is circularised.  There is a number of mechanisms proposed to explain anomalous sizes of inflated exoplanets (see \\eg Martin \\etal 2011 for references). Early studies suggested that a small but non-zero value of the orbital eccentricity of WASP-1~b, excited by an undetected planetary companion, could bloat the planet through tidal heating (Mardling 2007). This finding motivated us to study the WASP-1 system through the transit timing of WASP-1~b. ",
        "conclusions": "Our simultaneous observations of transits with two telescopes at different locations reduced mid-transit time uncertainties by 7--20\\%, depending on the quality of individual datasets. Our pilot observations without any filter, \\ie in white light, if performed in photometric conditions, result in higher data quality, and hence bring more precise mid-transit times.  The simultaneous analysis of our 13 new transit light curves and additional 6 from the literature results in redetermining system parameters which in most cases agree with the literature ones within $1\\sigma$. Despite the conservative approach, we derived more precise values than those ones published before.  The homogeneous analysis of transit light curves shows that transit times of the WASP-1~b planet follow the linear ephemeris, leaving no space for any sign of timing variation. Thus, we can only put observational constraints for a fictitious second planet in the system. The transiting planet has no detectable close planetary companion and this finding agrees with the analysis of a sample of hot Jupiter candidates observed with the {\\it Kepler} space telescope.     \\Acknow{We thank Mohammad Moualla for trying to obtain additional data. GM and DP acknowledge the financial support from the Polish Ministry of Science and Higher Education through the Iuventus Plus grant IP2011 031971. JN and MV would like to thank the project VEGA 2/0143/14 and the project APVV-0158-11. CM thanks the German national science foundation Deutsche Forschungsgemeinschaft (DFG) for support through the grant SCHR665/7-1. SR acknowledges the financial support from DFG in the grant NE 515/33-1. We acknowledge financial support from the Thuringian government (B 515-07010) for the STK CCD camera used in this project. Part of this paper is the result of the exchange and joint research project {\\em Spectral and photometric studies of variable stars} between Polish and Bulgarian Academies of Sciences. This research has made use of the Exoplanet Transit Database, maintained by Variable Star Section of Czech Astronomical Society.}"
    },
    "1402/1402.4806_arXiv.txt": {
        "abstract": "The hard X-ray continuum and gamma-ray lines from a Type Ia supernova dominate its integrated photon emissions and can provide unique diagnostics of the mass of the ejecta, the $^{56}$Ni yield and spatial distribution, its kinetic energy and expansion speed, and the mechanism of explosion. Such signatures and their time behavior ``X-ray\" the bulk debris field in direct fashion, and do not depend upon the ofttimes problematic and elaborate UV, optical, and near-infrared spectroscopy and radiative transfer that have informed the study of these events for decades.  However, to date no hard photons have ever been detected from a Type Ia supernova in explosion. With the advent of the supernova SN 2014J in M82, at a distance of $\\sim$3.5 Mpc, this situation may soon change.  Both NuSTAR and INTEGRAL have the potential to detect SN 2014J, and, if spectra and light curves can be measured, would usefully constrain the various explosion models published during the last $\\sim$thirty years.  In support of these observational campaigns, we provide predictions for the hard X-ray continuum and gamma-line emissions for fifteen Type Ia explosion models gleaned from the literature.  The model set, containing as it does deflagration, delayed detonation, merger detonation, pulsational delayed detonation, and sub-Chandrasekhar helium detonation models, collectively spans a wide range of properties and, hence, signatures. We provide a brief discussion of various diagnostics (with examples), but importantly make the  spectral and line results available electronically to aid in the interpretation of the anticipated data. ",
        "introduction": "\\label{intro}  A Type Ia supernova explosion is best modeled as a thermonuclear explosion (of carbon and oxygen, and perhaps some helium) of a degenerate object near one solar mass (M$_{\\odot}$), with a kinetic energy near 10$^{51}$ ergs and the production of $\\sim$0.5 to $\\sim$1.0 M$_{\\odot}$ of radioactive $^{56}$Ni (Pankey 1962; Colgate \\& McKee 1969; Clayton, Colgate, \\& Fishman 1969).  The object could be a Chandrasekhar white dwarf, the merger product of two white dwarfs \\citep{1993A&A...270..223K,1996ApJ...472L..81H}, or a slightly sub-Chandrasekhar white dwarf.  The explosions could be subsonic deflagrations \\citep{w7}, sub-Chandrasekhar detonations \\citep{1996ApJ...472L..81H}, delayed detonations \\citep{WW91,1992ApJ...393L..55Y,1995ApJ...443...89H,1998ApJ...495..617H}, or pulsating delayed detonations \\citep{1995ApJ...444..831H,1996ApJ...472L..81H} and could start in the deep interior, from distributed hot spots, or near the surface. The product is rich not only in radioactive $^{56}$Ni, but in intermediate-mass elements (such as Si, S, Ca, or Ar) and in non-radioactive iron-group isotopes.  Whatever the detailed mechanism and early (first $\\sim$1 second) explosive history, Type Ia supernovae have been studied for decades in the optical, ultraviolet, and near-infrared. Various systematics (such as the ``Phillips\" relation; Phillips 1993) have been uncovered and profitably used (Riess et al. 1998; Perlmutter et al. 1999), a range of peak brightnesses, atomic abundances, and $^{56}$Ni yields have been inferred, and many detailed models of the explosion itself have been proffered.  However, though a large fraction of the luminous energy of a Type Ia emerges in a hard X-ray continuum and gamma-ray lines, no Type Ia supernova has yet been observed and scrutinized at these energies.  While the optical spectra and light curves require non-LTE (LTE: local-thermodynamic-equilibrium) and time-dependent radiative transfer of thousands of observed energy levels, depending collectively upon hundreds of thousands to millions of atomic transitions for scores of ionization states (see, e.g., Pauldrach, Hoffmann, \\& Hultzsch 2014), the gamma-ray signatures of Type Ia models are relatively simple to generate and understand.  The energies and branching ratios of the finite forest of gamma-ray lines produced in the $^{56}$Ni($\\tau = 8.8$ days)$\\rightarrow$$^{56}$Co($\\tau = 111.3$ days)$\\rightarrow$$^{56}$Fe decay sequence are well understood, as are the mean lifetimes ($\\tau$) and the positron production factors.  The cross sections needed for transfer calculations of the gamma-ray lines are simply those of (inelastic) Compton scattering, photoelectric absorption, and (below $\\sim$70 keV) bremsstrahlung of the secondary electrons.  All models early in their expansion become homologous, freezing their density and abundance distributions. Therefore, the gamma-ray lines and hard-photon continuum fluxes as a function of time are ideal signatures of the geometry, abundances, abundance distributions, kinetic energies, and masses of the explosion. In a very real sense, they ``X-ray\" the debris, and provide the most direct fundamental constraints on Type Ia models.   Recently, a Type Ia supernova, now designated SN 2014J, exploded in the galaxy Messier 82 at a distance of $\\sim$3.5 megaparsecs (Zheng et al. 2014; Goobar et al. 2014; Cao et al. 2014; Kuulers et al. 2014). Previous theory (e.g., Diehl \\& Timmes 1998; Ambwani \\& Sutherland 1988; Burrows 1990; Burrows \\& The 1990; Burrows, Shankar, \\& Van Riper 1991; Bussard, Burrows, \\& The 1989; Chan \\& Lingenfelter 1991; H\\\"oflich et al. 1998a; Maeda et al. 2012; Milne et al. 2004; Summa et al. 2013) indicates that at such a close distance its gamma-ray lines with energies of 812, 847, and 1238 keV might be detectable by INTEGRAL/SPI (Roques, Schanne, \\& von Kienlin 2003; Vedrenne et al. 2003). For a 10$^6$-second integration, INTEGRAL/SPI has an approximate narrow-line 3-$\\sigma$ sensitivity at 1 MeV of $3.1\\times 10^{-5}$ photons cm$^{-2}$ s$^{-1}$. It also has a spectral resolution near 1.3 MeV of $\\sim$2.5 keV, sufficient in principle to resolve the various gamma-ray lines whose widths have been estimated to be 10$-$20 keV (Burrows \\& The 1990; Bussard, Burrows, \\& The 1989).  Even more, the supernova's hard X-ray Compton continuum from $\\sim$30 keV to $\\sim$80 keV should easily be observed by NuSTAR, boasting a $10^6$-second sensitivity of a few $\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ in this energy range (Harrison et al. 2013; Koglin et al. 2005).  The hard continuum of SN 2014J might also be observed by INTEGRAL/IBIS/ISGRI (Lebrun et al. 2003), with a 3-$\\sigma$ sensitivity from $\\sim$20 to $\\sim$100 keV of a few $\\times 10^{-6}$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ \\footnote{INTEGRAL/IBIS also has capabilities from $\\sim$15 keV to $\\sim$1 MeV, with the potential for detections at the $\\sim$few $\\times 10^{-5}$ cm$^{-2}$ s$^{-1}$ keV$^{-1}$ level between $\\sim$100 and 1000 keV (Isern et al. 2013). The Suzaku HXD detector (Takahashi et al.2007) extends to $\\sim$200 keV (and with reduced effective area to $\\sim$600 keV) and might have a shot at detecting this supernova at early times (perhaps near $\\sim$20 days after explosion) and at the higher continuum energies.}. Excitingly, and by the end of 2014, ASTRO-H will be launched and will have capabilities in the region below 100 keV comparable to those of NuSTAR.  Specfically, below $\\sim$100 keV, HXI on ASTRO-H will have a 3-$\\sigma$ sensitivity for a $10^6$-second integration of a few $\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ (Kokubun et al. 2010). SGD on ASTRO-H has a design sensitivity over the 100$-$450 keV range of 5-10$\\times 10^{-8}$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ for a $10^6$-second integration (Tajima et al. 2010).  There already exist in the literature many predictions for the hard photon signatures of many disparate Type Ia models (see above). The gamma-ray line light curves, line shapes, hard X-ray continuum flux spectrum, flux band ratios, and the ratio of the hard X-ray fluxes to those of the gamma-ray lines are together sensitive to and diagnostic of the $^{56}$Ni yields, $^{56}$Ni and high-Z element spatial distributions and mixing, total masses, explosion kinetic energies, explosion asymmetries, and the distinction between supersonic detonation and subsonic deflagration (for a review, see Burrows 1990). Depending upon model, the peak gamma-line flux occurs $\\sim$60$-$90 days after explosion, the Comptonization continuum from $\\sim$30 keV to $\\sim$1 MeV peaks earlier (near day 20$-$30), and is usefully diagnostic of model before day 60, and the spectral peak of the continuum is found near $\\sim$100$-$200 keV before day 100.  The fluxes below $\\sim$100 keV are very senstive to photoelectric absorption in the debris, which in turn is related to iron-peak yield and its spatial distribution.  If there are iron-peak isotopes near the debris surface, the fluxes below $\\sim$100 keV can be severely photoelectrically suppressed, while the corresponding gamma-ray lines fluxes near $\\sim$1 MeV are higher.  The Compton continuum flux above $\\sim$100 keV would also be higher.  Hence, the ratios of the hard X-ray flux below $\\sim$100 keV to the continuum flux above $\\sim$100 keV or to the gamma line fluxes near $\\sim$1 MeV are strongly dependent on the iron-peak distribution.  Correspondingly, if the $^{56}$Ni is more deeply buried in the debris, the hard X-ray fluxes below $\\sim$100 keV are enhanced and these systematics are reversed.  Moreover, if the explosion energies are large, and/or the $^{56}$Ni yields high, the $^{56}$Co-decay line fluxes at 847 and 1238 keV can be 50$-$100\\% higher. The line shapes directly reflect the distribution of the produced $^{56}$Ni in velocity space (Burrows \\& The 1990).  In fact, an analytic model for these line shapes, including Doppler shift effects and distinguishing the matter and $^{56}$Ni distributions, has been developed (Bussard, Burrows, \\& The 1989).  We surmise, however, that a focussed set of new calculations in aid of the ongoing SN 2014J campaign to detect its hard continuum and line emissions might be of use.  In this spirit, we calculate in this paper various theoretical hard-photon signatures for a representative set of Type Ia explosion models, put at the 3.5 megaparsec distance of M82. The results can be scaled to any other distance using the inverse square law. The model results are also provided in tabular form (at \\url{http://www.astro.princeton.edu/\\~{}burrows/sn2014j/}) to help observers engaged in measurement and interpretation efforts navigate the theoretical terrain. In this paper, we don't provide gamma-ray line profiles, but encourage the interested reader to explore the associated predictions of Burrows \\& The (1990) and the analytic model of Bussard, Burrows, \\& The (1989). ",
        "conclusions": "\\label{conclusion}  We have calculated the line fluxes from $^{56}$Ni and $^{56}$Co decay and escape for the dominant lines of this beta decay chain, the concomitant hard X-ray Compton and bremsstrahlung continua, and the iron fluorescence line fluxes for fifteen representative Type Ia explosion models. The inclusion of bremsstrahlung (for the first time for most of the fifteen models) is particularly important for the proper interpretation of the anticipated NuSTAR data. We have done this without preconceived notions of which explosion models or modalities might be preferred from previous considerations.  From the line light curves and ratios, continuum band fluxes and their temporal evolution, and rise times and peak magnitudes, one may hope to constrain the various physical parameters of the explosion, and, perhaps, eliminate classes of explosion models while determining explosion parameters. In principle, the $^{56}$Ni yields and distributions, kinetic energies, and total ejecta masses (among other things) can be measured.  Our model results are available in electronic form at \\url{http://www.astro.princeton.edu/\\~{}burrows/sn2014j/} to aid in the interpretation of the hard photon data anticipated in the ongoing SN 2014J, as well as future, observing campaigns."
    },
    "1402/1402.4117.txt": {
        "abstract": "We present a tally of Milky Way late-type dwarf stars in \\Nfields~ WFC3 pure-parallel fields (227 arcmin$^2$) from the Brightest of Reionizing Galaxies (BoRG) survey for high-redshift galaxies. Using spectroscopically identified M-dwarfs in two public surveys, the CANDELS and the ERS mosaics, we identify a morphological selection criterion using the half-light radius ($r_{50}$), a near-infrared J-H, G-J color region where M-dwarfs are found, and a V-J relation with M-dwarf subtype.  We apply this morphological selection of stellar objects, color-color selection of M-dwarfs and optical-near-infrared color subtyping to compile a catalog of \\nM\\ M-dwarfs belonging to the disk of the Milky Way with a limiting magnitude of $m_{F125W} < 24 (AB)$. Based on the M-dwarfs statistics, we conclude that (a) the previously identified North/South discrepancy in M-dwarf numbers persists in our sample; there are more M-dwarfs in the Northern fields on average than in Southern ones, (b) the Milky Way's single disk scale-height for M-dwarfs is 0.3-4 kpc, depending on sub-type, (c) the scale-height depends on M-dwarf subtype with early type (M0-4) high scale-height ($z_0 =3-4$ kpc) and later types M5 and above in the thin disk ($z_0 = 0.3-0.5$ kpc). (d) a second component is visible in the vertical distribution, with a different, much higher scale-height in the Southern Fields compared to the Northern ones. We report the M-dwarf component of the Sagittarius stream in one of our fields with 11 confirmed M-dwarfs, 7 of which are at the stream's distance % In addition to the M-dwarf catalog, we report the discovery of \\nT\\ T-dwarf and \\nL\\ L-dwarfs from their near-infrared colors. The dwarf scale-height and the relative low incidence in our fields of L- and T-dwarfs in these fields makes it unlikely that these stars will be interlopers in great numbers in color-selected samples go high-redshift galaxies. The relative ubiquity of M-dwarfs however will make them ideal tracers of Galactic Halo substructure with EUCLID and reference stars for JWST observations. ",
        "introduction": "% Adapted from Ryan 05 Counting stars to infer the shape and size of our Milky Way Galaxy is a classic experiment in Astronomy. However, it is also among the most prone to insufficient data as well as conceptual shortfalls \\citep[e.g., Herschel 1785;][]{Kapteyn22}. Over time, it was established that the Milky Way's scale height is inversely proportional to the masses of stars sampled \\citep{Gilmore83,Gilmore84,Siegel02}.  Initially, many of the Galactic models were focused on relatively luminous giant stars and rarely addressed (sub)stellar objects due to completeness issues \\citep[see for reviews of star counts and Galactic structure;][]{Bahcall86,Gilmore89, MWasGal, Majewski93}.  Interest has now shifted to the spatial distribution of (sub)stellar objects, in part because they are possibly the most significant interlopers in studies of extremely high redshift objects \\citep[see e.g.,][]{Caballero08}, but also because these are physically some of the longest lived objects. Consequently, any deep high-redshift survey, especially in the near-infrared, undertaken with the Hubble Space Telescope automatically becomes a census of the smallest Milky Way stars. % More importantly, these objects constitute the most numerous and oldest population of stars in the Galactic disk. % For instance, \\cite{Pirzkal05} determined the scale-height of different types of dwarfs from the Hubble Ultra Deep Field \\citep[HUDF,][]{Beckwith06}. \\cite{Ryan05} found L and T dwarfs in a small set of ACS parallel observations. \\cite{Stanway08} and \\cite{Pirzkal09} determined the Galactic scale-height of M-dwarfs from the Great Observatories Origins Deep Survey fields \\citep[GOODS,][]{goods}. \\cite{Ryan11} added a search of early WFC3 pure-parallel fields for L and T dwarfs. These studies gradually improved statistics on L, T and M-dwarfs to several dozens of objects.  The identification of these Galactic dwarfs in images is difficult as one needs both high-quality data to distinguish from extragalactic sources and a good understanding of the initial mass function or the local density of these objects as a class. Two of the above studies \\citep{Pirzkal05,Pirzkal09} benefited from spectroscopic identification of the dwarf subtype in the HST/ACS grism spectra of the HUDF and GOODS fields (the GRAPES\\footnote{Grism ACS Program for Extragalactic Science, see \\cite{Pirzkal04}} and PEARS\\footnote{Probing Evolution And Reionization Spectroscopically, see \\cite{Pirzkal09}.} projects respectively). Combined with an index of stellarity in the direct image, this made identification of the dwarfs near certain. Alternatively, \\cite{Kilic05} used high-precision proper motion to identify the Galactic objects. % Once an object has been identified as a candidate dwarf, its luminosity distance must be estimated, which requires accurate (sub)typing. Without direct spectra, we will show that this can be readily achieved with sufficient color coverage from the optical to the near-infrared. % A dwarf census including subtype and distance, can answer several outstanding issues regarding the shape of the Milky Way traced by its most numerous stellar. members.  \\cite{Pirzkal09} found, like \\cite{Stanway08}, an odd discrepancy in the numbers of M-dwarfs: there are 24\\% more M-dwarfs in the Northern than in the  Southern GOODS field. Since our vantage point in the Milky Way is above the plane of the disk, one would expect the Southern GOODS fields to have more M-dwarfs. However, as present, this discrepancy has only been determined from two sight-lines out of the Galaxy, so it is still possible this may be source count field-to-field variance.  By studying the older stellar population content of our Galaxy along many sight-lines down to the faintest possible magnitude, we can directly observe and measure the shape of the thin Galactic disk as traced by the old stellar population.  In order to improve our understanding of the contribution of faint dwarfs to Galactic structure, we will need much improved statistics to compare against Milky Way structure models. We use the recent {\\em HST/WFC3} Pure Parallel searches for bright high-redshift galaxies to find Milky Way dwarfs and determine the parameters of their distribution in the Milky Way disk \\citep{Trenti11,Yan10b, Bradley12}. Thus far, based on {\\em HST} imaging, the disk scale height for dwarfs has been found to be: 300 pc (Table \\ref{t:z0}). \\citep{Zheng01}, 400 $\\pm$ 100 pc \\citep{Pirzkal05}, 350 $\\pm$ 50 pc, \\citep[L, T dwarfs][]{Ryan05}, 370 $\\pm$ 65 pc (for M4--M9 dwarfs) and 300 $\\pm$ 70 pc \\citep[for M0--M9 dwarfs,][]{Pirzkal09} and 300 $\\pm$ 25 ($\\pm$ 31 systematic) pc. \\citep[$>$M8, ][]{Ryan11}.  The goal of this paper is to explore the reality of the North-South difference in M-dwarf counts and to accurately determine the scale-height of our Milky Way as traced by M-dwarfs. % Previously, \\cite{Ryan11} investigated the L and T dwarf content, here we focus our analysis on the M dwarfs. % If the reader is suddenly overwhelmed by a strong sense of deja-vu, \\cite{Ryan11} tackles a good number of our fields in a very similar manner but they look for L and T dwarfs, ignoring the bounty of M-dwarfs in the sample. We focus on the latter.  The paper is organized as follows, \\S \\ref{s:data} reports the properties of the {\\em HST/WFC3} observations we use, \\S \\ref{s:cat} discusses how we generate the relevant {\\sc sextractor} catalogs, \\S \\ref{s:cat:candels} presents our observational calibration of morphological identification of stars, M-dwarfs specifically, and subsequent sub-typing of the M-dwarfs. \\S \\ref{s:MWstars} presents our catalog of dwarfs in the BoRG fields and briefly discusses the validity of our morphological selection. In \\S \\ref{s:disc} we discuss the North/South difference, the Galactic scale-height as a function of dwarf subtype. \\S \\ref{s:concl} lists our conclusions and outlines future work.    % We propose to use the multi-filter WFC3 data from HST Proposal 11700, a pure parallel SNAP proposal, to put a limit the oblateness of the M-dwarf MW disk. % We have a reasonable idea already of the thin disk height from HUDF and GOODS-N/S data (GRAPES and PEARS grism projects respectively).  \\begin{table}[htdp] \\caption{\\label{t:z0}The scale-heights for the thin disk of dwarf stars from the literature.} \\begin{center} \\begin{tabular}{l l l} Milky Way Scale  \t& Reference\t& type stars \\\\ Height ($z_0$) (pc.)\t\t\t\t\t\t\t& & \\\\ \\hline 300  \t\t\t\t\t\t\t& \\cite{Zheng01} \t& \\\\ 400 $\\pm$ 100  \t\t\t\t& \\cite{Pirzkal05} \t& \\\\ 350 $\\pm$ 50 \t\t\t\t\t& \\cite{Ryan05} \t& L, T dwarfs\\\\ 370 $\\pm$ 65 \t\t \t\t\t& \\cite{Pirzkal09} \t& M4--M9 dwarfs\\\\ 300 $\\pm$ 70  \t\t\t\t\t& \\cite{Pirzkal09}\t& M0--M9 dwarfs \\\\ 300 $\\pm$ 25 $\\pm$ 31 \t\t\t& \\cite{Ryan11} & $>$M8 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}% ",
        "conclusions": "\\label{s:concl}   % Is there still a North/South difference? What is the scale-height in M-dwarfs? Does it depend on the subtypes?  \\begin{itemize} \\item Stars can be selected best using {\\sc sextractor}'s half-light radius ($r_{50}$), corroborating previous identifications (Figures \\ref{f:r50} and \\ref{f:borg:r50}). \\item M-dwarfs can be reliably sub-typed using a near-infrared and optical color (Figure \\ref{f:typing}). \\item Relatively more early-type M-dwarfs are identified in Northern fields as opposed to Southern Galactic ones. This confirms the dichotomy \\cite{Pirzkal09} noted (Table \\ref{t:ns}). \\item We report an over-density in one of our fields, borg\\_1230+0750, consistent with the M-dwarf members of the Sagittarius stream at a distance of 20-25 kpc. (Figure \\ref{f:sag:hist:d}). \\item A naive, single-component fit of the vertical distribution of M-dwarfs shows a steady decline of scale-height, $z_0$, with M-dwarf sub-type (Figure \\ref{f:Mtype-z0}). \\item The grouping of spurs of stars at different Galactic coordinates strongly hints at substructure in the Milky Way disk/halo of M-dwarfs (Figure \\ref{f:Mdwarf:z-MWcor}). \\end{itemize}"
    },
    "1402/1402.4989_arXiv.txt": {
        "abstract": "Polar crown prominences are made of chromospheric plasma partially circling the Sun's poles between 60\\degree\\ and 70\\degree\\ latitude. We aim to diagnose the 3D dynamics of a polar crown prominence using high cadence EUV images from the \\textit{Solar Dynamics Observatory} (SDO)/AIA at 304 and 171 \\AA\\ and the \\textit{Ahead} spacecraft of the \\textit{Solar Terrestrial Relations Observatory} (STEREO-A)/EUVI at 195 \\AA. Using time series across specific structures we compare flows across the disk in 195 \\AA\\ with the prominence dynamics seen on the limb. The densest prominence material forms vertical columns which are separated by many tens of Mm and connected by dynamic bridges of plasma that are clearly visible in 304/171 \\AA\\ two-color images. We also observe intermittent but repetitious flows with velocity 15 km s$^{-1}$ in the prominence that appear to be associated with EUV bright points on the solar disk. The boundary between the prominence and the overlying cavity appears as a sharp edge. We discuss the structure of the coronal cavity seen both above and around the prominence. SDO/HMI and GONG magnetograms are used to infer the underlying magnetic topology. The evolution and structure of the prominence with respect to the magnetic field seems to agree with the filament linkage model. ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=\\linewidth]{images/3cc.pdf} \\caption{a) Prominence/filament system observed on 13 February 2011 (a) STEREO-A 195 \\AA\\ b) SDO (304, 171, 193 \\AA\\ composite). Arrows point to common bright points in the STEREO and SDO images. Grid lines are SDO longitude and latitude separated by 15\\degree. The righthand longitude line shows the position of SDO limb. The lowest latitude line is 45\\degree. The STEREO image has been derotated to align with disk features at 00:05 UT on 15 February. A black box outlines the region presented in Figure~\\ref{footpoints}.} \\label{over} \\end{figure*}    As their name suggests, polar crown prominence/filament systems create a crown that approximately outlines the polar coronal holes \\cite{hir85,tand95}. They form between dispersed active region flux and polar coronal holes. Their position and number changes with the solar cycle \\cite{Webb84,McIntosh92}. Starting a few years before solar maximum, the phase in which the reported observations were made, polar crown filaments become more common \\cite{zir97,Minar98}. On the disk they appear as rows of dark pillars or vertical columns with intermittent filamentary connections \\cite{lin2003,sch10,dud12,li2013}. On the limb, persistent small-scale flows are seen in the cold plasma pillars \\cite{cha08,Berger08,Berger11} that often reside in long-lived coronal cavities \\cite{gib06,gib10,hab10,don11}.   One key to understanding the formation of polar crown prominences is the linkage of flux across and between dispersing active regions \\cite{Martens01,mac01}. Most filaments are found between initially unrelated active regions \\cite{gai97,Martin98,Mackay08} along the magnetic polarity inversion lines (PILs) which separate the regions of opposite polarity magnetic fields \\cite{Martin73}. It is probable that magnetic reconnection, driven by flows converging to the PIL, plays an important role in the filament formation and maintaining the filament mass (\\textit{e.g.} \\opencite{cha03}). As new connections are made between the adjacent active regions, open flux regions develop along the outer edges of the original bipoles \\cite{mac06}.  Prominences and their on-disk counterparts, filaments, have been thoroughly investigated with imaging and spectroscopic techniques (\\inlinecite{lab10} and references therein). They are relatively cool and dense structures suspended in the hot million-degree kelvin corona \\cite{hir85,Mackay10,lab10}. When observed in chromospheric lines, they appear to consist of two main structures: the spine that runs along the filament channel, and the barbs that project sideways from the spine and connect to the photosphere \\cite{Martin98,aul98}. Often in images of polar crown filaments only the barbs, sometimes known as the filament feet, are visible as dark pillars \\cite{sch10,dud12,li2013}. The two main reasons for prominences appearing dark in EUV images is absorption and/or volume blocking \\cite{anz05}. Absorption is due to the photoionization of H and He, by the coronal radiation passing through the prominence \\cite{kuc98,hei01,hei03}. Volume blocking is caused by the absence of coronal plasma in a volume of space which does not participate in the emission of coronal lines \\cite{sch04,hei08}. Some of the prominence features can also be visible as emission in the 171 \\AA\\ filter \\cite{par12}.  Prominences observed in \\Halpha\\ and \\HeII\\ have different morphological features \\cite{wangh98}. In observations, the \\Halpha\\ emission coincides with the 193\\AA\\ absorption because the optical thickness of the prominence observed in the \\FeXII\\ line is similar to the optical thickness in the \\Halpha\\ \\cite{anz05}. Spectroscopic observations of lines formed at transition region and coronal temperatures show plasma streaming through the corona with velocities up to 70 km s$^{-1}$ for typically 13 min \\cite{kuc03}. Such flows would therefore travel a distance of 50 000 km, equivalent to a supergranule cell diameter. One of the outstanding problems is to disentangle structures along the line-of-sight and to see the connections between the cool \\Halpha~and 304 \\AA\\ emission and the hot coronal emission.  An important and probably necessary condition for the stability of the prominence/filament system is an overlying coronal arcade above the cold prominence plasma and related coronal cavity \\cite{low95,Hudson99,Gibson06,hei08}. The arcade is nearly potential field and surrounds both the cavity and the prominence. Most models for the cavity and prominence invoke either a flux-rope \\cite{kup74,Priest89,Balle89,low95,aul98,low02} or a sheared-arcade \\cite{ant94,dev00,aul02}. In these models the dense prominence plasma lies in the dips of the field and the cavity is filled with the surrounding, low-density plasma. Typically the transverse fields of quiescent polar crown prominences are observed to have an inverse polarity with respect to their overlying arcade \\cite{leroy84,anz94,ant94,kup96}. This is consistent with both the flux-rope and sheared arcade models. Coronal magnetic field measurements favour a flux-rope configuration in at least one of the observed polar crown cavities \\cite{rach13,bak13}.   The cavity in extreme ultraviolet (EUV) images is sometimes seen on top of cooler prominence material with horns of cooler material outlining the base of the cavity \\cite{Regnier11,schmit13,panesar13}, and sometimes surrounding the prominence \\cite{berger12}. From the top, the cavity shows up as a dim region in EUV images \\cite{vas09}.   In the present work, we focus on the dynamics of a small polar crown prominence filament  system seen on the north-west limb from Earth using high cadence \\textit{Solar Dynamics Observatory} (SDO) images from the \\textit{Atmospheric Imaging Assembly} (AIA) \\cite{lem12} and near the central meridian in 195 \\AA\\ \\textit{Solar Terrestrial Relations Observatory Ahead} (STEREO-A) images from the \\textit{Extreme UltraViolet Imager} (EUVI) \\cite{Howard08}. The prominence is typical of the many small ones seen at high latitudes. Hazy 171 \\AA\\ emission extends from the top and sides of the main 304 \\AA\\ prominence pillars, into and from which plasma streams along spidery legs. Recently, \\inlinecite{su12} have also analysed SDO and STEREO quadrature observations to investigate the magnetic structure of a large polar crown prominence/filament system, in the days preceding its partial eruption. Here we use the stereoscopic perspective of SDO and STEREO-A to separate the filament into pillars and interconnecting flows. In particular we are able to use combined 304 and 171 \\AA\\ AIA images to track the flows through the corona, and to find the starting and end points of the flows on the disk. Using the two instruments in quadrature, we are also able to investigate the events leading to the appearance of EUV cavities around and over the prominence. To identify the magnetic configurations underlying the prominence, we have used line-of-sight magnetograms obtained by the \\textit{Helioseismic and Magnetic Imager} (HMI) \\cite{sch12} and the \\textit{Global Oscillation Network Group} (GONG) \\cite{harvey96}.  The plan of the paper is as follows: In Section 2, we  describe the observations, and  give details of the data analysis; In Section 3, we first give an overview of the polar crown prominence and the structure of the filament channel. Then we investigate the evolution of the prominence over about 40 hours so as to identify the main prominence pillars in the STEREO-A images. Finally in Section 4 we present a sketch of the prominence/filament system and relate it to the large-scale magnetic field configuration.  \\begin{figure*}[ht] \\centering \\includegraphics[width=0.9\\linewidth]{images/171feb1.pdf} \\caption{Evolution of the filament system from 5--9 February 2011: Sequences of SDO/AIA 171 \\AA\\ intensity images showing the growth of pillar F2 and  bright points (BPs) in the filament channel. The white box roughly outlines the region of the filament channel observed on the limb on 13 and 14 February 2011. The dark lane is indicated with an arrow in (c).} \\label{evolution} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.9\\linewidth]{images/193feb.pdf} \\caption{Evolution of the filament system from 5--9 February 2011: Sequences of SDO/AIA 193 \\AA\\ intensity images showing the pillar F2 and bright points (BPs) in the filament channel. The images are taken at the same time as those in Figure~\\ref{evolution}.} \\label{evo} \\end{figure*}  \\begin{figure*}[ht] \\centering \\includegraphics[width=0.9\\linewidth]{images/7feb-hmi.pdf} \\caption{Magnetic field structure on 7 February 2011: (a) SDO/HMI 6173 \\AA\\ averaged magnetogram; (b) SDO/AIA 171 \\AA\\ intensity image overlaid with contours of $\\pm $~10 G.} \\label{hmi} \\end{figure*} ",
        "conclusions": "%   \\begin{figure} \\centering \\includegraphics[width=0.7\\linewidth]{images/sketch1.pdf} \\caption{Sketch of our observations. Stars show sites of the various bright points. Pillars are marked as F1, F2 and F3. Positive and negative sign indicate magnetic polarity.} \\label{skt} \\end{figure}   We have investigated the structure and dynamics of a polar crown prominence observed in the northern hemisphere on 13--14 February 2011 seen simultaneously at the limb from SDO and at disk center from STEREO-A. The observations were made during the rising phase of the solar cycle when there was no obvious major polar coronal hole in either the SDO or the STEREO images of the northern pole. We have identified three main pillars (vertical columns) of prominence plasma and four bright points which produced persistent brigtenings. Using time-series of STEREO on-disk images we deduce that the pillars are separated by hundreds of arcsecs and connected by flows along the filament channel. The character of the start and end sites are determined using running difference time-slice images of on-disk emission which show distinctly different behaviour for bright points, flows, and pillars. During the two days studied, we noticed that microflaring (at the footpoint of pillar F1) produced 171 \\AA\\ prominence flows that rise up and stop abruptly at a sharp concave edge in the corona because the plasma is constrained to flow along the flux rope.   One of the motivations for this study was to observe and study the dynamics of a prominence-cavity system. Initially, the cavity was visible as a concave, bowl-shaped dark region in 171 \\AA\\ emission  above the prominence. It was particularly noticeable when there were flows between pillars F1 and F2. When the connection between the two main pillars, F1 and F2, broke,  the magnetic field in the prominence/filament system reconfigured and the prominence increased in height.  The dark cavity region seen in SDO coincided with a dark region close to the filament channel, in the STEREO disk images. Unlike regions of the channel with dense, prominence plasma, this darker region did not show signs of flows in the STEREO time-series. We therefore predict that the darkness was due to its low density, rather than cool prominence plasma and that it is the disk signature of the prominence cavity. We found a similar region of reduced 193 \\AA\\ emission at approximately the same position in the SDO images when the prominence/filament channel was on the disk from SDO. In both the STEREO and the disk SDO images, persistent brightening was seen along the edge of this region of reduced 195/193 \\AA\\ emission.   The presence of the cavity also became apparent after a small flare-like brightening and EUV wave inside the filament channel at 01:00 UT on 14 February. This caused a coronal dimming that was immediately visible in the STEREO images and then later as a bowl-shaped dark region in 171 \\AA\\ emission (e.g. \\opencite{gib06}) in the limb EUV images. It remained as a visible cavity in 171 \\AA\\ until the region rotated around the limb.   We therefore find that the presence of the cavity was made visible in the EUV images in two circumstances. The first was flows along the prominence on 13 February that highlighted the boundary between the prominence and the cavity and the second was due to heating and compression caused by an eruption-like event close to the filament channel on 14 February.  To understand the prominence/filament structures in relation to the large-scale magnetic field, we compared the sites of bright points and pillars to the underlying photospheric flux. For this purpose we used line-of-sight magnetograms from GONG and composite SDO and STEREO-A 195 \\AA\\ EUV images. Figure~\\ref{carr_euv_lob} shows EUV images and contours of the images overlaid on the GONG magnetic fields for the 13 February 2011. The interpretation of these images at such high latitudes is ambiguous but the large scale is still discernible. To the south west was strong positive polarity active region flux that has diffused to this position during the previous 27 days (Figure~\\ref{carr_euv_lob}a,b). A close-up of the diffused positive flux and higher latitude negative flux is shown in the HMI image (Figure~\\ref{hmi}a). The cancellation of the newly diffused positive flux with pre-existing negative flux elements is the reason for the many microflares at B1, B3 and B4 (Figure~\\ref{carr_euv_lob}).   A sketch of the inferred magnetic topology and sites of the prominence features is shown in Figure~\\ref{skt}. Since the observations were made in the rising phase of the solar cycle when the negative polarity is dominant over the northern pole, we have placed a region of negative field over the pole. The newly diffused active region is on the south west. The positions of the other magnetic flux regions are more speculative and are chosen to reflect the structure seen in the GONG map and the observation that filaments form between adjacent active regions. \\inlinecite{Martens01} proposed a flux linkage model to explain the chains of polar crown filaments. Here we have the situation where the arrival of new active region flux is merging with an already formed chain due to magnetic field reconnection at B2, B1, and the footpoint of pillar F1 which led to the injection of flux and plasma into the prominence.   In this paper we have demonstrated the possibility of being able to dissect polar crown prominences into individual pillars, and filaments flows using observations from spacecraft separated by about 90\\degree. These dynamics suggest the presence of complex magnetic geometries in the filament channel. The arrival of diffused active region flux seems to play an important role in the evolution of this prominence/filament system.    \\begin{acks} We would like to thank the referee for his/her constructive comments. SDO data are courtesy of the NASA/SDO AIA and HMI science teams. STEREO data are courtesy of the STEREO Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) team. GONG magnetogram data obtained by the Global Oscillation Network Group (GONG) project, managed by the National Solar Observatory, which is operated by AURA, Inc. under a cooperative agreement with the National Science Foundation. NKP acknowledges the facilities provided by the MPS. One of the authors, SKT, is supported by an appointment to the NASA Postdoctoral Program at the NASA Marshall Space Flight Center, administered by Oak Ridge Associated Universities through a contract with NASA. \\end{acks}"
    },
    "1402/1402.5031_arXiv.txt": {
        "abstract": "The distinction between modified gravity and quintessence or dynamical dark energy is difficult. Many models of modified gravity are equivalent to models of coupled quintessence by virtue of variable transformations. This makes an observational differentiation between modified gravity and dark energy very hard. For example, the additional scalar degree of freedom in $f(R)$-gravity or non-local gravity can be interpreted as the cosmon of quintessence. Nevertheless, modified gravity can shed light on questions of interpretation, naturalness and simplicity. We present a simple model where gravity is modified by a field dependent Planck mass. It leads to a universe with a cold and slow beginning. This cosmology can be continued to the infinite past such that no big bang singularity occurs. All observables can be described equivalently in a hot big bang picture with inflation and early dark energy. ",
        "introduction": "\\label{Introduction}  Einstein's equation  \\begin{equation}\\label{1} M^2(R_{\\mu\\nu}-\\frac{1}{2}Rg_{\\mu\\nu})=T_{\\mu\\nu} \\end{equation} expresses geometrical quantities on the left hand side in terms of matter and radiation on the right hand side. The basic geometrical quantity is the metric $g_{\\mu\\nu}$, with $R_{\\mu\\nu}$ and $R$ the Ricci tensor and curvature scalar formed from the metric and its derivatives. The energy momentum tensor $T_{\\mu\\nu}$ contains contributions from the particles of the standard model (``baryons'', neutrinos, radiation) and from dark matter.  The observation of the present accelerated expansion \\cite{SN1,SN2} as well as indications for an inflationary epoch in very early cosmology tell us that equation \\eqref{1} cannot be complete despite the numerous successful predictions of general relativity. One may supplement terms on the left or right side, as indicated by the dots \\begin{equation}\\label{2} R_{\\mu\\nu}-\\frac{1}{2}Rg_{\\mu\\nu}+_{\\dots}=\\frac{1}{M^2}(T_{\\mu\\nu}+_{\\dots}~). \\end{equation} Additional contributions to the energy momentum tensor are usually called dark energy, whereas a change on the left hand side is associated with a modification of gravity or general relativity.  It is obvious that this distinction cannot be a particularly strict one since the validity of an equation does not depend on where one writes terms. The most prominent candidate for the explanation of an accelerated expansion, the cosmological constant, can be interpreted as an additional contribution to the energy momentum tensor $\\Delta T_{\\mu\\nu}=\\lambda g_{\\mu\\nu}$. This interpretation is suggested by the contribution of the effective potential of the Higgs scalar to $\\lambda$, or similar for other scalar fields. We could write the cosmological constant term also on the left hand side and consider it as a modification of gravity - after all it influences the gravitational equations in ``empty space''.  One may try a more concise definition of the meaning of modified gravity by requiring that the change of the Einstein tensor on the l.h.s. of eq. \\eqref{1} involves derivatives of the metric, while terms with additional fields and no derivatives of $g_{\\mu\\nu}$ would contribute to $T_{\\mu\\nu}$. We will see, however, that modified gravity models defined in this way can often be rewritten in terms of different fields, frequently additional scalar fields. What appears in one field basis as a modification of gravity with terms involving derivatives of the metric shows up as dark energy with new fields and without metric derivatives in an other field basis. In particular, modified gravity theories that are consistent with the observed evolution of the universe are often equivalent to dynamical dark energy or quintessence. The borderline between modified gravity and dark energy becomes rather fuzzy. In fact, the first model of quintessence has originally been formulated as a modification of gravity \\cite{CW3}.  The reason for this ambiguity between modified gravity and dark energy is connected to a basic property: observables depend on the dynamical degrees of freedom, but not on the choice of fields used to describe them (``field relativity''). For example, the metric may contain a scalar degree of freedom besides the graviton. This scalar is not distinguished from a ``fundamental scalar field'' (cosmon) which is the basic ingredient of quintessence.  These lecture notes will present several examples for the equivalence of modified gravity and quintessence. In particular, $f(R)$ gravity or a large class of non-local gravity models are equivalent to coupled quintessence \\cite{CW2,ACQ}. We do not aim, however, to cover all possible modifications of gravity. More general modified gravity models may contain further non-scalar degrees of freedom (vectors of tensors), involve an infinite number of degrees of freedom, or give up the basic diffeomorphism symmetry underlying general relativity.  Recent reviews of modified gravity can be found in refs. \\cite{ReVO,Rev1,Rev2,Rev3,Rev3A,ReV4}. We concentrate here on the deep connection between modified gravity and coupled quintessence. This helps to understand many of the rich features of modified gravity in a simple and unified way. It also shows that many claims for observational distinguishability between modified gravity and quintessence are actually not justified.  In sects. \\ref{Basic setting} and \\ref{Weyl scaling} we display our basic setting and discuss the field transformations that relate different versions of a given physical model. In sect. \\ref{Brans-Dicke cosmology} we describe the cosmology of Brans-Dicke theory in the language of coupled quintessence. This points to strong observational bounds on the effective coupling $\\beta$ between the cosmon and matter that will play an important role later. Sect. \\ref{Scalar tensor models} discusses general scalar-tensor models with actions containing up to two derivatives. We highlight the importance of field-dependent particle masses in order to find models obeying the bounds on $\\beta$. Sect. \\ref{Slow Freeze Universe} discusses a simple three-parameter cosmological model along these lines which is compatible with all present observations from inflation to late dark energy domination. Formulated as a scalar-tensor theory (Jordan frame) it exhibits an unusual cosmic history. The universe shrinks during the radiation- and matter-dominated epochs and the evolution is always very slow. Cosmological solutions remain regular in the infinite past and there is no big bang singularity. On the other hand, the same model is characterized in the Einstein frame by a more usual big bang picture. This underlines that the field transformations that a crucial for these notes also incorporate important conceptual aspects.  In sect. \\ref{Modified gravity with $f(R)$} we describe the equivalence of $f(R)$-modified gravity with coupled quintessence \\cite{CNO,HS,KAS}. For constant particle masses the equivalent coupled quintessence models exhibit a large universal cosmon-matter coupling $\\beta=1/\\sqrt{6}$. This issue is a major problem for the construction of realistic $f(R)$ models. We sketch in sect. \\ref{$f(R)$-gravity with varying particle masses} how a vanishing coupling $\\beta=0$ can be obtained for $f(R)$-models with field-dependent particle masses. In sect. \\ref{Non-local gravity} we turn to simple models of non-local gravity. Again, such models are equivalent to coupled quintessence. In sect. \\ref{Higher derivative} we ask the general question to what extent modified gravity models which lead to second order field equations, as Horndeski's models \\cite{Horn}, can find an equivalent description as coupled quintessence models. We find a huge class of such modified gravity models for which the scalar-gravity part is given by the action for quintessence, while additional information is contained in the details of the effective cosmon-matter coupling. Our conclusions are drawn in sect. \\ref{Conclusions}. Parts of sects. \\ref{Scalar tensor models}, \\ref{Slow Freeze Universe} have overlap with work reported in refs. \\cite{CWVG,Wetterich:2014eaa}. ",
        "conclusions": "\\label{Conclusions}  Can one distinguish modified gravity from dark energy by observation? In view of the equivalence of a large class of modified gravity models with coupled quintessence an answer to this question is not straightforward. Statements that modified gravity and quintessence lead to different growth factors for cosmic structures apply only to quintessence models without coupling to matter. We have seen, however, that the quintessence models that are equivalent to modified gravity typically have a nonzero coupling $\\beta$ between the cosmon and different forms of matter. (This coupling needs not to be the same for all species of massive particles.) In this view precision measurements of the growth rate can differentiate between uncoupled and coupled quintessence and determine bounds on $\\beta$. The issue if there are modified gravity models that can be distinguished observationally from coupled quintessence is much harder to answer.  Modified gravity models almost always involve new degrees of freedom besides the graviton. This is a consequence of the fact that models for a massless spin two particle are severely constrained by consistency requirements. The conjecture that consistency requires diffeomorphism symmetry (more precisely its unimodular subgroup) has never been proven, but no counter examples are known either. A model containing a massless spin two particle as the only degree of freedom is then rather close to general relativity. Modifications of gravity therefore typically involve additional degrees of freedom, as scalars or massive spin two particles.  The field description of the additional degrees of freedom is not unique. For example, a scalar may be described as a component of the metric (modified gravity) or by a separate field (quintessence). Very large classes of models can be mapped onto each other by non-linear field transformations. Field relativity states that observables cannot depend on the choice of fields. For models related by field transformations no observational distinction is possible. We have seen that this holds for variable gravity models where the Planck mass is field dependent. It also applies to $f(R)$-models and large classes of non-local gravity. Very general models equivalent to coupled quintessence models have been discussed in the preceding section.  For all these models modified gravity and coupled quintessence should merely be seen as two different pictures describing the same reality, in analogy to the Jordan frame and Einstein frame for the metric. For practical computations of the evolution of homogeneous cosmology and fluctuations around this background the simplest way uses the Einstein frame. This holds both for the linear treatment of fluctuations and for numerical simulations in the non-linear regime. The physical effects of the cosmon-matter coupling $\\beta$ are intuitively accessible in the Einstein frame.  For modified gravity models that are equivalent to coupled quintessence one may ask: why then discuss them all? If there is no observational distinction, the discussion of such modifications of gravity may at first sight look like a redundant exercise. A deeper answer concerns questions of simplicity and naturalness. Models of modified gravity can be very simple and involve no unnatural parameters. Nevertheless, the equivalent description in the Einstein frame by coupled quintessence may hide simplicity and naturalness in the complexity of the field transformation. An example is the big bang singularity. We have presented in sect. \\ref{Slow Freeze Universe} a modified gravity model for which the ``beginning'' of the universe is very slow and cold. It has no big bang singularity, the cosmological solution can be continued to the infinite past $t\\to -\\infty$. In the Einstein frame the same model is described as a hot big bang. Models may be regular in the Jordan frame and show a big bang singularity in the Einstein frame. This singularity is then due to a singularity in the field transformation  \\cite{CWU}, in close analogy to a coordinate singularity.  The question of naturalness is often closely linked to symmetries. Scale symmetry is explicitly visible in the modified gravity description of the models in sects. \\ref{Scalar tensor models}, \\ref{Slow Freeze Universe}. It is realized by a multiplicative rescaling of the metric and the scalar field $\\chi$. In the presence of quantum fluctuations scale symmetry is violated by $\\chi$-dependent (``running'') dimensionless couplings. For fixed points of the running exact (quantum-) scale symmetry is restored. For the quantum effective action \\eqref{SF2} such fixed points are present for $\\chi\\to 0$ and $\\chi\\to\\infty$ \\cite{Wetterich:2014eaa}.  In our model in sect. \\ref{Slow Freeze Universe} the asymptotic value \\be\\label{CCA} \\lambda_\\infty=\\lim_{\\chi\\to\\infty} V(\\chi)/\\chi^4 \\ee vanishes for the fixed point at $\\chi\\to\\infty$. This can be motivated by properties of a possible ultraviolet fixed point in dilaton quantum gravity \\cite{RP} or by dilatation symmetry in higher dimensions \\cite{CWDL1,CWDL2}. The fixed point with $\\lambda_\\infty=0$ is the deeper reason why the cosmological constant vanishes asymptotically in the Einstein frame, $\\lim_{\\varphi\\to\\infty}V'(\\varphi)\\to 0$. Without this understanding of naturalness as a consequence of fixed point properties one would argue in the Einstein frame that naturalness suggests the addition of a constant to eq. \\eqref{F6}. Apparently convincing qualitative arguments on the induction of a cosmological constant by quantum fluctuations in the Einstein frame yield very different results when applied in the Jordan frame. A constant term in $V(\\chi)$ yields a term $V'(\\varphi)\\sim \\exp (-2\\alpha\\varphi/M)$ in the Einstein frame which vanishes for $\\varphi\\to\\infty$. This is one more example how modified gravity can shed new light on questions of naturalness.  The possibility of field transformations from modified gravity theories to coupled quintessence models in the Einstein frame is an extremely useful tool for the discussion of observational consequences of a model. It should not prevent us, however, to look for modified gravity theories distinguished by simplicity and naturalness.                 \\newpage"
    },
    "1402/1402.1534_arXiv.txt": {
        "abstract": "Astrophysical cosmology constrains the variation of Newton's Constant in a manner complementary to laboratory experiments, such as the celebrated lunar laser ranging campaign. Supernova cosmology is an example of the former and has attained campaign status, following planning by a Dark Energy Task Force in 2005. In this paper we employ the full SNIa dataset to the end of 2013 to set a limit on G variation. In our approach we adopt the standard candle delineation of the redshift distance relation. We set an upper limit on its rate of change   $|\\dot{G}/G|$ of 0.1 parts per billion per year over 9 Gyrs. By contrast lunar laser ranging tests variation of G over the last few decades. Conversely, one may adopt the laboratory result as a prior and constrain the effect of variable G in dark energy equation of state experiments to $\\delta$w~$<$~0.02. We also examine the parameterization G $\\sim$ 1+z. Its short expansion age conflicts with the measured values of the expansion rate and the density in a flat Universe. In conclusion, supernova cosmology complements other experiments in limiting G variation. An important caveat is that it rests on the assumption that the same mass of $^{56}$Ni is burned to create the standard candle regardless of redshift. These two quantities, f and G, where f is the Chandrasekhar mass fraction burned, are degenerate. Constraining f variation alone requires more understanding of the SNIa mechanism. ",
        "introduction": "\\label{sec:intro}   The Planck mass, $m_P = \\surd \\hbar c/G$, is a fundamental quantity of stellar astrophysics. The Chandrasekhar mass to order unity is $m_C ~=~m_P^3/m_p^2$, where $m_p$ is the proton mass. The maximum mass of a star is approximately % the Chandrasekhar mass multiplied by the square of the ratio of radiation pressure to gas pressure (Eddington 1917). The minimum mass of a black hole is within order unity of the Planck mass. Variation of Newton's constant affects supernova cosmology via change in the Planck mass over cosmic time and was first considered by Gazta\\~{n}aga \\etal (2002). At that time there were 42 SNe available; there are now 581 (Suzuki \\etal 2012).  Uzan (2003) and Narimani, Moss, \\& Scott (2010, 2012) advise that constraining the constancy of dimensional quantities is perilous. Preferred quantities are, for example, the ``gravitational fine structure constant\",  $\\alpha_{\\rm g}\\equiv G m_{\\rm p}^2/\\hbar c$. The lunar laser ranging experiment initiated by NASA's Apollo mission is an $\\alpha_{\\rm g}$ experiment, monitoring the specific potential energy of the Earth-Moon system $Gm_{\\rm p}^2/ct$, where t is the time of flight of Earth launched photons. Measuring the luminosity distance of galaxies, $D_L$, from type Ia supernovae  is an $\\alpha_{\\rm g}$ experiment, as $D_L^2 ~\\propto~ m_C/m_p$, assuming a fixed fraction, f, of $m_C$ is turned into energy and stellar luminosities are calibrated by hydrogen burning stars. Specifically, to within a numerical constant of order unity, $m_C/m_p~=~\\alpha_{\\rm g}^{-3/2}$.  Speculation about varying G began with Dirac (1937), who noted that the ratio of the electrostatic and gravitational forces between an electron and a proton was of the same order as the number of times an electron orbits a proton in the age of the Universe. He conjectured that $\\alpha_{\\rm g}$ might decay as the inverse of cosmic time. This 20th century gravity problem (which is sometimes tackled anthropically) has been totally eclipsed in the last decade by the cosmological constant problem (Sol\\`{a} 2013). The contribution to the vacuum energy density of fluctuations in the gravitational field % is larger than is observationally allowed by some 120 orders of magnitude. Instead, the vacuum energy density is of the same order of magnitude as the present mass density of the universe. Although ongoing type Ia supernova observations indicate that the equation of state of Einstein's General Relativity is the best fit, this gross cosmological constant problem provides no comfort for constant G  orthodoxy.  Garcia-Berro \\etal (2007) review astronomical measurements and constraints on the variability of fundamental constants generally. Garcia-Berro \\etal (2006) fit a polynomial  G(z) = $G_o(1-0.01z+0.3z^2-0.17z^3)$ to the SNIa data, suggestive of a G larger in the past. Verbiest \\etal (2008) measure orbital period rates of pulsars and set a limit of $|\\dot{G}/G| = 23 \\times 10^{-12} yr^{-1}$. From white dwarf cooling, Garcia-Berro et al. (2011) derive an upper bound  $\\dot{G}/G = -1.8 \\times 10^{-12} yr^{-1}$ and Corsico \\etal (2013) find a white dwarf pulsation limit of $\\dot{G}/G = -1.3 \\times 10^{-10} yr^{-1}$. Tomaschitz (2010) considers a  gravitational constant scaling linearly with the Hubble parameter, and fits the SNIa Hubble diagram and AGN source counts, concluding that further observational constraints are required.  Furthermore, the luminosity of degenerate carbon core supernovae is proportional to the mass of carbon burned to $^{56}$Ni. The precise mechanism which powers a type Ia supernova explosion is a matter of lively debate, and we do not know yet whether a detonation or a deflagration occurs. What it is clear is the close correlation between the mass of nickel synthesised in the outburst and the luminosity. This is discussed in quantitative experiments by Gazta\\~{n}aga \\etal (2002). The type SNIa standard candle is thought to result from a high degree of regulation of this quantity, such as would be provided by approach to a physical limit, the Chandrasekhar mass. However, one may conjecture that the fraction of $m_C$ that is turned into energy may also vary with z. Like variation of G with z, this is also an issue for constraints on the equation of state of the Universe arising from supernova measurements. ",
        "conclusions": "Our conclusions from this work are as follows. \\begin{enumerate} \\item The validity of the SNIa standard candle depends on the stability of G and the stability of f, the fraction of the Chandrasekhar mass turned into energy. We have considered the former in this paper and derived a constraint on the gravitational fine structure constant which can be compared with other combined astrophysical-cosmological constraints. But this is inextricably degenerate with possible evolution of f due to changes over cosmic time of SNIa progenitor astrophysics. With this caveat we set a SNIa cosmology constraint on G stability  in the standard cosmology of (--3,+7.3) $\\times$ 10$^{-11}$ per year.  \\item This limit is two orders of magnitude weaker than that from lunar laser ranging. But that is a current epoch result and complements, but does not replace, a constraint that spans cosmic time.  \\item The limit is also two orders of magnitude weaker than that arising from the great sensitivity to density of Big Bang nucleosynthesis. The SNIa standard candle, however, has the distinct advantage of covering the last 10$^{10}$ years of cosmic time, rather than the first twenty minutes.  \\item Our result is an update of Gazta\\~{n}aga \\etal (2002), who found $\\dot{G}/G~<$ 12 $\\times$ 10$^{-12}~ h_{70}$ /yr for $\\Omega_\\Lambda$ = 0.8, $\\Omega_M$ = 0.2. This is a 1$\\sigma$ limit, like ours, and directly comparable since $h_{70}~\\approx$ 1. The order of magnitude more supernovae now available have allowed us to relax their flat Universe assumption, but has also relaxed their limit on $\\dot{G}$.  \\item Caution would dictate that experiments to measure the equation of state of the Universe carry the caveat that f and G stability is assumed. For the latter, lunar laser ranging is available as a prior and limits $\\delta w$ to 0.02, but the former has not been quantified and demands further understanding of the SNIa mechanism.  \\item A parameterization G $\\sim$ 1+z is interesting on theoretical grounds related to the unity of forces. However, with such an equation of state the expansion age of the Universe is too short. That parameterization can therefore be rejected with 99\\% confidence.   \\end{enumerate}"
    },
    "1402/1402.0638_arXiv.txt": {
        "abstract": "{The Cherenkov Telescope Array (CTA) is the next generation of Imaging Atmospheric Cherenkov Telescopes. It would reach unprecedented sensitivity and energy resolution in very-high-energy gamma-ray astronomy. In order to reach these goals, the systematic uncertainties derived from the varying atmospheric conditions shall be reduced to the minimum. Different instruments may help to account for these uncertainties. Several groups in the CTA consortium are currently building Raman LIDARs to be installed at the CTA sites. Raman LIDARs are devices composed of a powerful laser that shoots into the atmosphere, a collector that gathers the backscattered light from molecules and aerosols, a photosensor, an optical module that spectrally select wavelengths of interest, and a read-out system. Raman LIDARs can reduce the systematic uncertainties in the reconstruction of the gamma-ray energies down to 5\\% level.  All Raman LIDARs subject of this work, have design features that make them different than typical Raman LIDARs used in atmospheric science, and are characterized by large collecting mirrors ($\\sim2~$m$^2$). They have multiple elastic and Raman read-out channels (at least 4) and custom-made optics design. In this paper, the motivation for Raman LIDARs, the design and the status of advance of these technologies are described.} ",
        "introduction": "Currently in its design stage, the Cherenkov Telescope Array (CTA) is an advanced facility for ground-based very-high-energy  gamma-ray astronomy~\\cite{bib:actis}. It is an international initiative to build the next-generation Cherenkov telescope array covering the energy range from a few tens of GeV to a few hundreds of TeV with an unprecedented sensitivity. The design of CTA is based on currently available technologies and builds upon the success of the present generation of ground-based Cherenkov telescope arrays (H.E.S.S., MAGIC and VERITAS\\,\\footnote{\\url{www.mpi-hd.mpg.de/hfm/HESS/},\\\\ \\url{wwwmagic. mppmu.mpg.de}, \\\\\\url{veritas.sao.arizona.edu}}).  Nowadays, the main contribution to the systematic uncertainties of imaging Cherenkov telescopes stems from the uncertainty in the height- and wavelength-dependent atmospheric transmission for a given run of data. Atmospheric quality affects the measured Cherenkov yield in several ways: the air-shower development itself, the loss of photons due to scattering and absorption of Cherenkov light out of the camera field-of-view, resulting in dimmer images and the scattering of photons into the camera, resulting in blurred images. Despite the fact that several supplementary instruments are currently used to measure the atmospheric transparency, their data are only used to retain good-quality observation time slots, and only a minor effort has been made to routinely correct data with atmospheric information~\\cite{bib:nolan1,bib:dorner,bib:reyes}. This situation may change with the CTA atmospheric monitoring and calibration program~\\cite{Doro:2013swa}. There are several goals behind this program. The first is to increase the precision and accuracy in the energy and flux reconstruction throughout the use of one or more atmospheric instruments. Secondly, a precise and continuous monitoring of the atmosphere will allow for an increase of the telescope duty cycle with the extension of the observation time during hazy atmospheric conditions, which are normally discarded in the current experiments because of the uncertainty in the data reconstruction. Finally, a possible ``smart scheduling'', i.e., an adaptation -- if required -- of the observation strategy during the night that considers the actual atmospheric condition can be activated with a precise monitoring program of the atmospheric conditions.  In this document, we mainly focus on the importance of the use of Raman LIDARs to retrieve the differential atmospheric absorption at different wavelengths at the observatory site. We will show why the Raman LIDAR is expected to be an optimal instruments for CTA. There are currently three groups developing independently Raman LIDARs for CTA, but the interest around these instruments is increasing and possibly other groups will try to develop different solutions. In this document, we present the main design ideas behind these LIDARs and the progress in their construction.  \\vspace{-3mm} ",
        "conclusions": "Monte Carlo simulations showed that while data affected by enhancements of the ground layer can be scaled rather easily up to high levels of extinction, this is not the case anymore for (cirrus) clouds at altitudes from 6 to 12~km a.s.l., which create strong energy-dependent effects on the scaling factors. Moreover, the images from atmospheric showers are distorted depending on the location of the shower maximum, which varies even for showers of a same energy. Very high altitude layers, in turn, produce only effects on the very low energy gamma-ray showers.  Depending on the properties of the chosen site for CTA, still to be decided, it would probably make sense to create $10-20$ typical atmospheric simulations within these possibilities and interpolate between them. The natural solution is the use of (Raman) LIDARs, which were described in this contribution. However, the use of complementary instruments that measure integral or differential (in altitude) atmospheric parameters is possible and envisaged. Once retrieved the differential atmospheric transparency, different strategies may be foreseen to accurately and precisely reconstruct data, ultimately reducing the reconstructed energy and flux uncertainties. This can be achieved by re-calibrating the data themselves, either event-wise or bin-wise, or by simulating adapted atmospheres~\\cite{Doro:2013swa}. In addition, it would be possible to increase the duty cycle of the telescopes by retrieving those data taken during non-optimal atmospheric conditions which are normally discarded by standard clean-atmosphere analysis, especially important during e.g. multi-wavelength campaigns or target of opportunity observations.  Three independent groups are currently developing custom-made Raman LIDARs. In Argentina, CEILAP is developing a novel-design 6-mirror Raman LIDAR, while in Europe, IFAE/UAB in Spain and LUPM in France are adapting a telescope from the terminated CLUE experiment as Raman LIDAR. Status and plans were reported in this proceedings.     \\vspace{3mm} \\footnotesize{{\\bf Acknowledgment:}{ We gratefully acknowledge support from the agencies and organizations listed in this page: \\url{http://www.cta-observatory.org/?q=node/22}}}"
    },
    "1402/1402.2298_arXiv.txt": {
        "abstract": "By depleting the ionizing photon budget available to expand cosmic HII regions, recombining systems (or Lyman limit systems) can have a large impact during (and following) cosmic reionization.   Unfortunately, directly resolving such structures in large-scale reionization simulations is computationally impractical.  Instead, here we implement a sub-grid prescription for tracking inhomogeneous recombinations in the intergalactic medium.  Building on previous work parameterizing photo-heating feedback on star-formation, we present large-scale, semi-numeric reionization simulations which self-consistently track the local (sub-grid) evolution of {\\it both} sources {\\it and} sinks of ionizing photons. Our simple, single-parameter model naturally results in both an extended reionization and a modest, slowly-evolving emissivity, consistent with observations. Recombinations are instrumental in slowing the growth of large HII regions, and damping the rapid rise of the ionizing background in the late stages of (and following) reionization.  As a result, typical HII regions are smaller by factors of $\\sim$2--3 throughout reionization. The large-scale ($k\\lsim0.2$ Mpc$^{-1}$) ionization power spectrum is suppressed by factors of $\\gsim2$--3 in the second half of reionization. Therefore properly modeling recombinations is important in interpreting virtually all reionization observables, including upcoming interferometry with the redshifted 21cm line.  Consistent with previous works, we find the clumping factor of ionized gas to be $C_{\\rm HII}\\sim4$ at the end of reionization. ",
        "introduction": "\\label{sec:intro}  Reionization is a story of sources and sinks of ionizing photons. Historically, the most famous protagonists of this tale have been the sources: the tiny ($\\lsim$ percent level) fraction of baryons which are able to condense and form stars inside the first galaxies. A fraction of the ionizing photons emitted from these stars (and accreting black holes) manage to escape their host galaxies and begin ionizing the surrounding intergalactic medium (IGM). Driven by the birth and evolution of galaxies, these cosmic HII regions grow and overlap, eventually permeating all of space and completing the last major phase change of our Universe: hydrogen reionization.  On the other hand, the ``sinks'' of ionizing photons have received comparably less attention. An ionizing photon escaping a galaxy can have one of two fates: (i) it can ionize a neutral atom beyond the surrounding HII region, contributing to its growth; or (ii) while passing through the surrounding HII region it can encounter and reionize an atom which was previously ionized but subsequently recombined.  These later sinks (often dominated by so-called Lyman limit systems, LLSs, though they can also include less dense, more diffuse structures) could substantially delay reionization by depleting the photon budget available for expanding the HII regions (e.g. \\citealt{MHR00, CSSF06, QLZD07, FOOD12, KG13}).  Furthermore, sinks can substantially alter reionization morphology, imposing a limit to the size of isolated HII regions (e.g. \\citealt{FO05, MMS12, AA12}).  Understanding their impact is crucial in interpreting all observations of the epoch of reionization.   LLSs are better understood at lower redshifts ($z\\lsim4$), post-reionization. Radiative transfer calculations \\citep{QOF11, RPRS13}, together with observations of quasar $\\text{Ly}\\alpha$ spectra (e.g. \\citealt{POW10, SC10}), confirm that the general properties of LLSs can be understood in terms of the local Jeans length \\citep{Schaye01}, although the detailed properties are difficult to simulate as they (and especially even more dense systems) can occur very near galaxies.  Most IGM recombinations are sourced by gas with densities near the critical threshold required for self-shielding (e.g. \\citealt{MHR00, FO05}).  Unfortunately, extending this simple picture of sinks to the epoch of reionization is not straightforward.  From a modeling standpoint, as always the main difficulty lies in the enormous dynamical range required. One needs to fully resolve the internal structure of the dominant population of self-shielded regions (with typical sizes of $\\sim$ ten proper kpc; e.g. \\citealt{Schaye01})\\footnote{If the IGM is cold prior to reionization, the Jeans length could even be 1--2 orders of magnitude smaller (e.g. \\citealt{ETA13}).} in cosmological simulations which sample reionization structure on scales of $\\gtrsim$ hundreds of comoving Mpc (e.g. \\citealt{FHZ04}).  Although single-box simulations of reionization have made significant progress lately (for example, see the recent review in \\citealt{TG11}), tackling this full range of scales remains computationally out of reach for the foreseeable future.  Moreover, the properties of the recombining gas are sensitive to the local density and radiation fields, both of which can be very inhomogeneous and can evolve rapidly during reionization.  Conclusions regarding the role of sinks must therefore be robust to astrophysical uncertainties in these quantities.    Here we implement a sub-grid model of recombinations inside large-scale, semi-numerical simulations of reionization.  Each cell's recombination history depends on the local values of density, ionizing UV background (UVB), and reionization history.  Building on \\citet{SM13b}, our simulations also include the effects of photo-heating on the star formation rate; therefore, we include reionization feedback on both the sources and the sinks.  We investigate the relative impact of inhomogeneous recombinations on the reionization history and morphology, as well as the evolution of the UVB, emissivity, mean free path to ionizing photons and the clumping factor (see below) of the ionized gas.  This paper is organized as follows. In Section \\ref{sec:reion} we describe our large-scale reionization simulations.  In Section \\ref{sec:subgrid} we present our sub-grid prescriptions to model UVB feedback on galaxies (\\S \\ref{sec:UVB_feedback}), and the small-scale structure of the IGM (\\S \\ref{sec:fluct}). In Section \\ref{sec:results} we present our results, investigating the role of sinks and sources on IGM properties during reionization.  In Section \\ref{sec:concl} we present our conclusions.  Throughout we assume a \ufb02at $\\Lambda$CDM cosmology with parameters ($\\Omega_{\\rm m}$, $\\Omega_{\\rm \\Lambda}$, $\\Omega_{\\rm b}$, $h$, $\\sigma_{\\rm 8}$, $n$) = ($0.28$, $0.72$, $0.046$, $0.70$, $0.82$, $0.96$), as measured by the Wilkinson Microwave Anisotropy Probe (WMAP; \\citealt{WMAP13}), and also consistent with recent results from the Planck satellite \\citep{Planck13}. Unless stated otherwise, we quote all quantities in comoving units. ",
        "conclusions": "\\label{sec:concl}  We study the role of ionizing photon sinks during reionization by implementing a sub-grid recipe inside large-scale, semi-numeric simulations of reionization.  Building on previous work, we self-consistently model the evolution of both sinks and sources of ionizing photons, using prescriptions calibrated to numerical simulations.  The inhomogeneous recombination rate and emissivity in our model depend on the local density, photo-ionizing background, and reionization history.  We find that both UVB feedback on sources, and recombinations slow the growth of large HII regions, prolonging reionization and resulting in a more uniform ionization morphology.  Although recombinations are more potent, both affects amplify one another as they most strongly affect the same biased regions.  Such regions were the first to ionize, allowing enough time for both recombinations and photoevaporation to take effect.  In our complete model, the end of reionization is delayed by $\\Delta z \\approx 2.5$.  This delay is sufficient for simple reionization models to match best estimates of the mid and end stages of reionization, as implied by recent observations of the CMB, quasar spectra, and LAEs.  Recombinations are mostly responsible for the slowing and eventual \"freeze-out\" of large HII regions ($\\gsim$10 Mpc).  This results in a dramatic suppression of the large-scale ($k\\lsim0.2$ Mpc$^{-1}$) ionization power-spectrum by factors of $\\gsim 2$--3, even as early as $\\avenf < 0.5$.  Such a dramatic impact on these scales makes recombinations invaluable in interpreting upcoming data from 21cm interferometers.  Furthermore, recombinations are responsible for damping the rapid rise of the mean free path and photo-ionizing background during the late stages of reionization.  The volume-averaged photoionization rate increases by a modest factor of $\\sim2$ (as opposed to $\\sim5$ ignoring recombinations) during the last, $\\avenf<0.2$ stages of reionization.  Our complete model naturally results in an early start, and \"photon-starved\" end of reionization, as well as a modest, slowly-evolving emissivity (governed by UVB feedback which depletes gas from increasingly massive halos as time progresses).  Although undoubtedly too simplistic, this physical picture ameliorates empirically-motivated claims for a rapid redshift evolution in galaxy properties, like the escape fraction of ionizing photons (e.g. \\citealt{HM12, KF12})\\footnote{To be more precise, our emissivity estimates are only sensitive to the product on the RHS in eq. (\\ref{eq:zeta}); most importantly $f_{\\rm esc} f_\\ast$, as the dependence on $f_{\\rm b}$ is accounted for in our UVB feedback treatment via an evolving $M_{\\rm crit}(\\textbf{x},z)$.  Our model assumes that this product does not have a halo mass or redshift dependence. In our model, the emissivity during reionization is dominated by faint galaxies, 1-2 orders of magnitude fainter than current detection limits.  Instead, at lower redshifts ($z\\lesssim3$--4), the emissivity from star-forming galaxies is likely dominated by the observed LBGs, which do seem to have somewhat lower values of $f_{\\rm esc} f_\\ast$ than our fiducial choices in eq. \\ref{eq:zeta} (see for example, Fig. 7 in \\citealt{KF12}). This is suggestive of at least a mild {\\it halo mass} dependence of $f_{\\rm esc} f_\\ast$ (e.g. \\citealt{AFT12}).}.  Further progress would be aided by improved, physically-motivated models for the evolution of the ionizing emissivity.  \\vspace{0.5cm}  We thank M. McQuinn, M. Haehnelt, J. Bolton and S. Furlanetto for insightful comments on draft versions of this work."
    },
    "1402/1402.5471_arXiv.txt": {
        "abstract": "{We present large scale observations of C$^{18}$O (1--0) towards four massive star forming regions for MON R2, S156, DR17/L906 and M17/M18. The transitions of H$_2$CO (1$_{10}$--1$_{11}$), C$^{18}$O (1--0) and 6 cm continuum were compared towards the four regions. Analysis of observation and Non--LTE model results shows that the brightness temperature of the formaldehyde absorption line is strongest in background continuum temperature range of about 3 -- 8 K. The excitation of the H$_2$CO absorption line is affected by strong background continuum emission. From the comparison of H$_2$CO and C$^{18}$O maps, we found that the extent of H$_2$CO absorption is broader than that of C$^{18}$O emission in the four regions. Except for the DR17 region, the H$_2$CO absorption maximum is located at the same position with the C$^{18}$O peak. The good correlation between intensities and widths of H$_2$CO absorption and C$^{18}$O emission lines indicate that the H$_2$CO absorption line can trace dense and warm regions of the molecular cloud. Finding that N(H$_2$CO) was well correlated with N(C$^{18}$O) in the four regions and that the average column density ratio is $<$N(H$_2$CO)/N(C$^{18}$O)$>$ $\\sim$ 0.03. ",
        "introduction": "% \\label{sect:intro}  Dense molecular cloud cores are considered to be sites of star formation. Studying the structure and condition of these cores allows us to have a better understanding of the whole process of star formation. It is important to have a comprehensive understanding of the evolution of dense regions.  The first detection of H$_2$CO was reported by Snyder et al. (1969). It is the first organic polyatomic molecule discovered in the interstellar medium. The dense clumps in the line of sight may be the region of H$_2$CO production (Prasad \\& Huntress 1980). The formaldehyde absorption provides a powerful tool for analyzing the distribution of dense interstellar gas (Scoville \\& Solomon 1972). For the typical densities of clouds near HII regions (10$^{3}$ to 10$^{5}$ cm$^{-3}$) the 1$_{10}$--1$_{11}$ transition of H$_2$CO is ideal (Gardner et al. 1984). The 4.83 GHz 1$_{10}$--1$_{11}$ transition of interstellar H$_2$CO is easily observed in absorption against galactic continuum source. And it is also easily seen in absorption toward many dark clouds where the molecule's  lines are absorbing the cosmic background radiation (CMB).  Downes et al. (1980) found that the intensity of the H$_2$CO absorption line is affected strongly by the background continuum emission in the HII region. Zhang et al. (2012) obtained a similar result in giant molecular clouds, and inferred that the contribution of CMB radiation to the H$_2$CO intensity was very weak. The giant HII regions have many compact continuum components, and the H$_2$CO absorption is affected by the distribution of these continuum components. A mapping study of 12 bright galactic HII regions and 2 dark clouds showed that the strong H$_2$CO absorption were quite away from the bright continuum peaks in a number of cases (Bieging et al. 1982). Pipenbrink \\& Wendker (1988) surveyed the brightest features in the Cyg X region and found that there are 22 positions where no H$_2$CO absorption line was detected. This is especially surprising in the high brightness levels of HII regions.  Recently, Tang et al. (2013) studied the correlation among H$_2$CO (1$_{10}$--1$_{11}$), $^{12}$CO (1--0), and $^{13}$CO (1--0) towards the MON R2, S156, DR17/L906 and M17/M18 regions, and they found that the H$_2$CO and $^{13}$CO integral intensity maps had analogous shapes, sizes, peak positions and molecular spectra presenting similar central velocities and line widths. Such a good agreement indicates that the H$_2$CO absorption and the $^{13}$CO emission lines arise from similar regions. The $^{12}$CO (1--0) line was not suitable for tracing the star formation region because of the larger optical depth. The $^{13}$CO (1--0) spectra tend to be saturated towards the high density regions, making it difficult to understand the detailed distribution of the high density regions. The abundance of C$^{18}$O molecule is less than that of $^{13}$CO molecule by a factor of $\\sim$ 5. The C$^{18}$O line is optically thin compared to $^{13}$CO and could trace dense regions ($>$ 10$^{4}$ cm$^{-3}$) (Dame et al. 1986). Therefore, the C$^{18}$O is a good spectral line to trace dense regions in nearby star formation regions. Minn et al. (1996) had compared the distributions of the H$_2$CO 6-cm absorption and C$^{18}$O (1--0) emission (Fuller et al. 1991) towards the dark cloud B5 region. They found that the two molecular lines showed a similar spatial distribution with both peaks at the same position. One thing H$_2$CO and C$^{18}$O have in common is that they both have been probed to trace the dense gas. The 1$_{10}$--1$_{11}$ transition of H$_2$CO is a unique probe to trace high-density gases at low temperature regions. Liszt \\& Lucas (1995) and Liszt et al. (2006) have found that the N(H$_2$CO) are quite well correlated with N(HCO${^+}$), N(C$_2$H), N(HCN), N(CS) and N(NH$_3$) in the diffuse clouds. Toward the cold, dense pre-protostellar cores, the H$_2$CO shows depletion (Young et al. 2004). Therefore, the H$_2$CO will show an interesting difference towards the warm, high-density regions.  In this paper, we report present comparison of H$_2$CO, C$^{18}$O and continuum on a scale of 2 $\\sim$ 10 pc towards four galactic HII regions of MON R2, S156, DR17/L906 and M17/M18. We are interested in two issues: (1) to seek the relation between the H$_2$CO line and the background continuum; (2) to make a comparative study on the H$_2$CO and C$^{18}$O lines. ",
        "conclusions": "\\subsection{Comparison between H$_2$CO and background continuum} The relation between H$_2$CO line temperatures and continuum temperatures have been noted by Zhang et al. (2012) in four HII regions. They concluded that the H$_2$CO integration intensity is positively correlated with the background continuum temperature when the continuum temperature is below 6 K. Those regions and the four regions Tang et al. (2013) observed have a low range of continuum temperature (T$_C$ $<$ 10 K), thus it is difficult to know the varied process of the H$_2$CO intensity in the high brightness continuum region. So we selected some sources (including giant molecular clouds (GMCs) and HII sources) from the previous observations of H$_2$CO and continuum, and plotted these sources together in Fig.1.  A horizontal slice in figure shows a U shape of the H$_2$CO intensity. The absorption line intensity is proportional to background continuum intensity as the continuum intensity is below about 10 K, which agrees with the conclusion of Zhang et al. (2012). Above this value, the intensity of H$_2$CO line shows slight dispersion, and may weakly correlate with the background continuum intensity. It suggests that the H$_2$CO intensity is strongly influenced by the strong background continuum emission. Actually, H$_2$CO absorption line intensities depend on the background continuum level, but they also depend on H$_2$CO column density and gas density (e.g., Garrison et al. 1975; Wootten et al. 1980).  \\begin{figure}[t] \\vspace*{2mm} \\begin{center} \\includegraphics[height=5.5cm]{ms1555fig1.eps} \\end{center} \\caption{The relation between the brightness temperature of H$_2$CO and background continuum.} \\end{figure}  We compared the Non--LTE of interstellar line spectra model (van der Tak et al. 2007) results with the observed results. The Non--LTE model shows similar distribution of the relation between the H$_2$CO line intensity and the background continuum intensity (see Fig.A.1). The absorption line intensity is proportional to the background continuum intensity, which guarantees a rough proportionality between 6 cm continuum and H$_2$CO intensity until the 6 cm continuum reaches some critical values corresponding to the background continuum brightness of about 3 -- 8 K. Above the critical value, the intensity of H$_2$CO line decreases as the background continuum temperature increases. Under this condition where the H$_2$CO cloud has a strong background continuum emission field, the H$_2$CO cloud can also be heated, so part of H$_2$CO cloud could be thermalized. The excitation temperature of H$_2$CO absorption line would increase above 2.7 K. This suggests that the strong background continuum emission affects the excitation of the H$_2$CO absorption line. It could be a reason why the strong H$_2$CO absorption were quite away from the bright continuum peaks in some clouds.  \\begin{figure}[t] \\vspace*{2mm} \\begin{center} \\includegraphics[width=10.2cm,angle=0]{ms1555fig2.eps} \\end{center} \\caption{Contour maps of integrated area toward (a) MON R2, (b) S156, (c) DR17/L906 and (d) M17/M18. The black contour, green contour and blue contour respectively indicates the integrated intensity of the H$_2$CO absorption line, H110$\\alpha$ recombination line and C$^{18}$O emission line. H$_2$CO and H110$\\alpha$ data were selected from Tang et al. (2013). (a) For the MON R2 region: C$^{18}$O contour levels are 0.20 to 0.66 in steps of 0.06 K km s$^{-1}$. (b) For the S156 region: C$^{18}$O contour levels are 0.20 to 0.46 in steps of 0.05 K km s$^{-1}$. (c) For the DR17/L906 region: C$^{18}$O contour levels are 0.22 to 0.49 in steps of 0.05 K km s$^{-1}$. (d) For the M17/M18 region: C$^{18}$O contour levels are 0.80 to 2.12 in steps of 0.27 K km s$^{-1}$. The gray bars are given in units of K km s$^{-1}$ for the negative integrated intensity of H$_2$CO.} \\end{figure}  The statistic of observed results shows that approximately 75\\% of observed positions with the line-to-continuum ratio $|$T$_L$/T$_C$$|$ are less than 0.2 and the average value $<$$|$T$_L$/T$_C$$|$$>$ is about 0.066 which is higher than the value 0.052 estimated by Whiteoak \\& Gardner (1974) for 280 galactic radio sources or source components. Towards the MON R2, S156 and L906 regions the $|$T$_L$/T$_C$$|$ ratios are higher than those towards the DR17 and M17 HII regions. Tang et al. (2013) detected strong H110$\\alpha$ and continuum emission in DR17 and M17 regions, but did not detect H110$\\alpha$ emissions towards MON R2, S156 and L906 regions which have weak continuum emission. It is suggested that the presence of strong H110$\\alpha$ and 6 cm continuum emissions is the reason why high $|$T$_L$/T$_C$$|$ ratios do not occur in these regions.  \\subsection{Comparison between H$_2$CO absorption and C$^{18}$O emission} We compared the velocities of H$_2$CO and C$^{18}$O, and there is a good agreement between two molecules towards four regions (see Table B.1). The distributions of H$_2$CO and C$^{18}$O towards four regions are displayed as a line integral intensity maps in Fig.2, which shows that H$_2$CO and C$^{18}$O have quite a similar geometry and peak positions. The H$_2$CO absorption maximum is located at the same position with C$^{18}$O peak except DR17 region. The map shows that the extent of the H$_2$CO absorption is broader and smaller for the C$^{18}$O emission in four regions. This indicates that the two molecules are probing different parts of the molecular clouds. The C$^{18}$O emission is more likely tracing denser gas than the H$_2$CO absorption is. 176 H$_2$CO absorption points were detected in four regions, but we only detected about 52 points with the C$^{18}$O emission. The C$^{18}$O is somewhat more difficult to be detected than the H$_2$CO is. Those maybe because the density of the H$_2$CO absorption line is lower than that required to excite the C$^{18}$O. The comparison between $^{12}$CO and $^{13}$CO mapping results towards four regions (Tang et al. 2013), shows that $^{12}$CO, $^{13}$CO, H$_2$CO and C$^{18}$O distribution ranges are gradually decreasing. The correlation between the H$_2$CO and the $^{13}$CO distributions is better than that between the H$_2$CO and $^{12}$CO, C$^{18}$O distributions.  \\begin{figure}[t] \\vspace*{2mm} \\begin{center} \\includegraphics[height=5.5cm]{ms1555fig3.eps} \\end{center} \\caption{Correlation between the H$_2$CO and C$^{18}$O lines fluxes. The dashed line is the linear fit for H$_2$CO and C$^{18}$O flux data.} \\end{figure}  \\begin{figure}[h] \\vspace*{2mm} \\begin{center} \\includegraphics[height=5.5cm]{ms1555fig4.eps} \\end{center} \\caption{Figure shows the line widths of H$_2$CO and C$^{18}$O. The dashed line is the linear fit for H$_2$CO and C$^{18}$O line widths, and the correlation coefficient is 0.7.} \\end{figure}  The relation of fluxes between H$_2$CO and C$^{18}$O was represented. We selected the fluxes data shown as squares in Fig.3, and used them to make linear fits. The best linear fit to a straight line is, {\\setlength\\arraycolsep{2pt} \\begin{eqnarray} \\log I(H_2CO)=(0.60\\pm0.06)\\log I(C^{18}O)-(0.38\\pm0.02) (K km s^{-1}). \\end{eqnarray} It shows that I(H$_2$CO) is linearly well correlated with I(C$^{18}$O), and the correlation coefficient is 0.8. The H$_2$CO absorption line detections depend on the background continuum brightness and the gas density. For the C$^{18}$O lines, their intensities are dominated by the kinetic temperature of the gas. The fitted relation between I(H$_2$CO) and I(C$^{18}$O) indicates that H$_2$CO absorption line can trace the dense and warm region of the molecular cloud.  The line widths show the thermal motion and turbulence of internal kinematical properties of molecular clouds. H$_2$CO and C$^{18}$O molecules have the same molecular weight, so the thermal line width is about 0.2 km s$^{-1}$ for two molecules, which is concluded at 30 K. All H$_2$CO and C$^{18}$O line widths observed at large scale exceed the thermal line width. In addition, the H$_2$CO line has hyperfine structure components. The contribution of hyperfine structure and thermal broadening to the measured H$_2$CO line widths is likely to be small to moderate (Tang et al. 2013). The correlation between H$_2$CO and $^{12}$CO, $^{13}$CO line widths are not obvious in these four regions (Tang et al. 2013). Towards the line widths between H$_2$CO and C$^{18}$O, it shows a good correlation (see Fig.4). And there is a similar line widths range of 1 -- 5 km s$^{-1}$ for two tracers. These indicate that the line broaden mechanism of H$_2$CO and C$^{18}$O lines could be similar in the dense regions.  \\begin{figure}[t] \\vspace*{2mm} \\begin{center} \\includegraphics[height=5.5cm]{ms1555fig5.eps} \\end{center} \\caption{Histogram shows the histogram of the peak optical depths of H$_2$CO and C$^{18}$O distribution.} \\end{figure}  \\begin{figure}[h] \\vspace*{2mm} \\begin{center} \\includegraphics[height=5.5cm]{ms1555fig6.eps} \\end{center} \\caption{Correlation of H$_2$CO and C$^{18}$O column density. The dashed line is the linear fit for H$_2$CO and C$^{18}$O column density, and the correlation coefficient is 0.7.} \\end{figure}  The derived peak optical depths of H$_2$CO and C$^{18}$O show that most of the H$_2$CO optical depths are slightly higher than those of C$^{18}$O (see Fig.5). The average optical depth of H$_2$CO is $<$$\\tau$(H$_2$CO)$>$ $\\sim$ 0.05 where the C$^{18}$O also have been detected. For C$^{18}$O, the average optical depth is $<$$\\tau$(C$^{18}$O)$>$ $\\sim$ 0.03. The similar average optical depths of C$^{18}$O and H$_2$CO indicate that the two tracers can both probe high density regions. The relation between H$_2$CO and C$^{18}$O peak column densities shows that the H$_2$CO has a rapid increase in abundance and is not as rapid as that of C$^{18}$O (see Fig.6). For the four regions, N(H$_2$CO) is well correlated with N(C$^{18}$O) and the average column density ratio is $<$N(H$_2$CO)/N(C$^{18}$O)$>$ $\\sim$ 0.03. The column density of H$_2$CO may be underestimated because the cloud size is smaller than our beam size. The true cloud size is overestimated (see e.g. Gardner \\& Whiteoak, 1972; Slysh,1975), therefore the column density N(H$_2$CO) may be higher than we estimated. At the same time, the column density of C$^{18}$O could be underestimated too."
    },
    "1402/1402.0853_arXiv.txt": {
        "abstract": "Cosmic rays, thermal gas and magnetic fields in FRII radio cavities are assumed to come entirely from winds flowing from just behind the jet shocks.  Combining analytic and computational methods, it is shown that the computed radio-electron energy distribution and synchrotron emissivity spectra everywhere in the Cygnus A radio cavity agrees with radio observations of the Cygnus A lobes.  The magnetic field energy density is small everywhere and evolves passively in the post-shock wind.  Most synchrotron emission arises in recent post-shock material as it flows back along the radio cavity wall.  Because it experienced less adiabatic expansion, the magnetic field in this young backflow is larger than elsewhere in the radio lobe, explaining the observed radio synchrotron limb-brightening. The boundary backflow decelerates due to small cavity pressure gradients, causing large-scale fields perpendicular to the backflow (and synchrotron emission) to grow exponentially unlike observations.  However, if the field is random on subgrid (sub-kpc) scales, the computed field reproduces both the magnitude and slowly decreasing radio synchrotron emissivity observed along the backflow. The radio synchrotron spectrum and image computed with a small-scale random field agree with VLA observations.  The total relativistic energy density in the post-jet shock region required in computations to inflate the radio cavity matches the energy density of relativistic electrons observed in the post-shock region of Cygnus A.  This indicates that the component in the jet and cavity that dominates the dynamical evolution is a relativistic pair plasma. ",
        "introduction": "Our recent 2D gas dynamical calculations of the remarkable FRII radio-X-ray source Cygnus A explicitly include the dynamical effects of both hot thermal gas and relativistic gas inside the radio cavity (Mathews \\& Guo 2010, 2012; MG10 and MG12 hereafter). Since the dynamics of relativistic, radio-synchrotron-emitting electrons are known, we now extend the results of MG12 by computing the radio synchrotron emission and magnetic flux throughout the radio cavity and compare with detailed VLA observations of Cygnus A.  Unlike all previous FRII computations, we do not explicitly compute the jet evolution, but instead assume that the energy that inflates the radio cavity enters via the well-observed post-jet shock region. It is assumed that all thermal, non-thermal and magnetic energy inside the radio cavity originates in this small region just behind the jet shock. This intensely energetic and possibly turbulent region is the source of a powerful wind that flows away in every direction except, we assume, upstream in the jet direction. With these assumptions, it is possible to compute the entire evolution of Cygnus A with a post-shock source region that moves out into the cluster gas; the jet is not explicitly computed. Avoiding the jets is convenient since sufficiently detailed observations of faint FRII jets do not currently exist while observations of the physical properties of the luminous post-shock regions are much better understood.  One of the chronic computational difficulties experienced by previous attempts to compute FRII dynamics, like those of Cygnus A, is the Kelvin-Helmholtz instability (KHI) that invariably occurs when gas, after leaving the jet post-shock region, is diverted to flow back in the anti-jet direction toward the cluster center as it inflates the radio cavity. The large velocity shear generated by this backflowing gas gives rise to KHIs that cause the computed cavity to become irregular in shape and to produce chaotic, irregular internal velocity fields. Neither of these attributes are observed in Cygnus A or are common in other FRII sources. Many or most FRII radio cavities have smoothly curved outer boundaries and observed radio-synchrotron ages that decrease in a smooth, monotonically fashion along the jet direction, showing no evidence that the interiors of radio cavities are scrambled by large scale KH-induced turbulence although turbulence on smaller scales is still possible. In particular, VLA observations of the radio cavity of Cygnus A reveal a very smooth, gently curved boundary that appears to be stable to KHI.  Computations driven by a moving post-shock source region allow us to reproduce a number of important radio and X-ray features observed inside the Cygnus A radio cavity. One of our most important results in MG12 is the realization that relativistic gas backflowing from the jet-shock is restricted to a rather narrow shell just adjacent to the outer boundary of the radio cavity. Since this flow carries with it the youngest, most energetic radio-emitting electrons and the largest fields, we conclude that radio synchrotron emission from the cavity arises mostly or exclusively from this boundary backflow. Although this radiating shell geometry has not been considered by most observers who interpret radio or inverse Compton X-ray emission from the radio lobe, edge-brightened radio synchrotron emission has indeed been verified in Cygnus A by Carvalho et al. (2005) and in other FRII sources by Daly et al. (2010).  The KHI in the boundary backflow must be damped. This can be accomplished either with a magnetic field along the flow direction or by adding a small viscosity. In MH12 we adopted viscous damping and this choice is further supported in the discussion in Section 2 below. Viscous damping arises naturally from the entrainment of a small amount of cluster gas into the rapid backflow. Entrainment also decelerates the low-inertia, thermally relativistic backflow, but most of the backflow deceleration is due to a small negative radial pressure gradient inside the cavity. (The gradient is small because of the large scale height of hot cavity gas in the cluster gravity field.) In turn, the deceleration significantly amplifies large scale magnetic fields in the backflow that are not parallel to the backflow velocity.  Observations of Cygnus A indicate that magnetic energy densities throughout the cavity region are less than the combined energy densities of local thermal and non-thermal gas components, and this is the assumption made in MG12. Consequently, the Lorentz term ${\\bf j \\times B}$ does not appear in the equation of motion, i.e. the large scale magnetic field is not strong enough to influence the flow of gas or relativistic plasma, and the frozen-in magnetic field evolves passively as it flows from the post-shock region into the radio cavity. This passivity allows us to explore the effects of various field morphologies in the boundary backflow using a post-processing procedure based on the same gasdynamical evolution described in MG12.  The radio-synchrotron emission throughout the radio cavity in Cygnus A is completely determined by the energy distribution of relativistic synchrotron electrons observed in the post-shock region $n(\\gamma)$ between energy $\\gamma = E/m_e c^2$ and $\\gamma + d\\gamma$. In Section 3 we describe how simple analytic solutions can be used to translate this post-shock energy spectrum to an evolved spectrum anywhere anywhere in the lobe and in particular along the boundary backflow.  Another new result discussed in MG12 is that the small regions of luminous radio and X-ray emission near the tips of the jets, often referred to as ``hotspots'', are more complex than previously recognized. These kpc-sized regions of intense radio synchrotron emission in FRII sources are not in general located at the energetic jet-shock source regions described above, but in arc-shaped regions just ahead along the jet direction where the post-shock wind first crashes against dense (shocked) cluster gas and is powerfully compressed. In Cygnus A radio and synchrotron-self Compton X-ray emission just behind the jet shock is much weaker than from the bright arc about 1.5 kpc ahead. Nevertheless, the post-shock region is certainly the hottest spot in the radio cavity, i.e. the hottest spot is not the observed ``hotspot''. Fortunately, in at least one very powerful FRII source it is possible to observe X-radiation from both spot regions (Erlund et al. 2010), so a new terminology is needed to distinguish between the post-shock and compressed arc-shaped regions. VLA observers describe significant radio structure in and near the hotspots (e.g. Black et al. 1992 and Leahy et al. 1997) and occasionally suggest a spatial distinction between the jet termination shock and bright radio emission, but it is unclear how these features relate to the spatial distinction we discuss here for which we provide a dynamical model.  In MG12 we chose to refer to the (typically fainter) post-shock region as the ``hotspot'' and the observationally brightest region as the ``offset''. Here we drop the use of ``hotspot'' altogether and unambiguously refer to the post-shock region as the ``shock spot'' and to the more luminous arc-shaped wind compression offset ahead as the ``bright spot''.  Detailed hydrodynamic modeling for Cygnus A such as we describe in MG12 and discuss here at length, requires considerable computing and analysis and therefore must be restricted to a small  set of well-chosen parameters such as the core-hotspot distance, the age, the magnetic field, jet luminosity, initial properties of the cluster gas in which the event occurred, etc. Among the best known parameters is the physical size of Cygnus A from core to hotspots, determined from its redshift with $H_0 = 75$ km s${-1}$ Mpc$^{-1}$, is about 60 kpc with 1 kpc per arcsecond (e.g. Wilson, Smith \\& Young 2006). We adopt an age of 10 Myrs for Cygnus A which can be determined from the spectral evolution of the radio continuum (e.g. Machalski, Chyzy, Stawarz \\& Koziel 2007) who find 10.4 $\\pm$ 1.6 Myrs using equipartition fields or from the size and strength of the bow shock and a simple self-similar evolutionary model in the cluster atmosphere ($\\sim 3$ Myrs from Wilson, Smith \\& Young 2006). Other previous estimates of the age do not differ significantly from our chosen value, 10 Myrs (e.g. Begelman \\& Cioffi 1989; Carilli et al. 1991; Kaiser \\* Alexander 1999). The average rate that cosmic ray energy is supplied to the radio cavity by the jet is approximately $L_{cr} = 10^{46}$ ergs s$^{-1}$ (e.g. Wilson, Smith \\& Young 2006). Our assumed value of $L_{cr}$ refers only to the single radio lobe we calculate, so the total power is $L_{tot} = 2L_{cr}$. However, this power is sufficient to inflate radio cavities to volumes only about half that observed in each Cygnus A cavity after 10 Myrs. Consequently, in \\S6 below we consider a second hydrodynamic model with $L_{cr} = 2.65 \\times 10^{46}$ ergs s$^{-1}$ that does reproduce the observed cavity volume. Magnetic field strengths measured by combining radio and X-ray observations are independent of equipartition assumptions and expected to be reasonably accurate, but they are computed with simpler cavity emission geometries than we describe here. Yaji et al. (2010) determine a (presumably uniform) field of 20 $\\mu$G in the Cygnus A radio cavity by comparing radio synchrotron emission with inverse Compton X-ray emission upscattered from the microwave background and synchrotron-self Compton (SSC) X-rays produced by upscattered synchrotron emission. Their analysis (especially for SSC) depends on the geometric morphology of the synchrotron-emitting region which Yaji et al. (2010) assume is uniform throughout the radio lobe. This differs considerably from the edge-brightened radio lobe emission we describe in MG12. Nevertheless, until a new data analysis is performed using our model, we adopt 20 $\\mu$G as a reference ``observed'' value for the radio lobe field. Our adopted field in the bright spot, about 220$\\mu$G, taken from Stawarz et al. (2007), is less dependent on global dynamical models. Overall, uncertainties in these key Cygnus A parameters, and the dynamical models we compute from them, are likely to be dominated more by systematic errors due to inadequate underlying assumptions than by statistical observational errors. ",
        "conclusions": "Observations of the Cygnus A radio cavity indicate that the magnetic field is nowhere strong enough to directly influence the gas plus CR dynamics, $u_B < e_c + e$. Nevertheless, the frozen-in field connects the dynamical motion of relativistic and non-relativistic components in our two-fluid hydrodynamical models of Cygnus A evolution (MG12). Because the magnetic field evolves passively with the post-shock hydrodynamic flow, a variety of different field evolution models can be investigated and compared with observation without re-computing the gasdynamics. The goal is to find transport properties, field morphologies, etc. that best fit detailed observations. By this means we hope to understand better the complex, poorly understood small scale physics in weakly magnetized plasmas dominated by relativistic energy density.  \\vskip.1in\\noindent$\\bullet$ A new terminology is proposed that avoids use of the standard ``hotspot'' designation near the tips of FRII jets. The ``shock spot'' is the immediate post-shock region that serves as the source of a wind that transports cosmic rays, gas and magnetic field into the Cygnus A radio cavity. The \"bright spot'' is the much more radio luminous region where the shock spot wind impacts against the dense (shocked) cluster gas just ahead.  \\vskip.1in\\noindent$\\bullet$ The Kelvin-Helmholtz instability (KHI) in FRII radio lobes must be damped to match the smooth appearance of observed radio cavities and the monotonic variation of observed synchrotron spectral ages inside the cavity. It is unlikely that the KHI in Cygnus A can be stabilized by large-scale magnetic fields along the direction of the boundary backflow inside the radio cavity. For stability the Alfv\\'en speed $v_A = B/(4 \\pi \\rho)^{1/2}$ must exceed the velocity difference across the shearing layer. To achieve this in Cygnus A either (i) the magnetic field in this one direction must be $\\sim 10$ times larger than the total field observed in the radio cavity or (ii) the density $\\rho$ of non-relativistic gas must be 100 times smaller than assumed in the MG12 calculation. But the ultra-high backflow velocity at these ultra-low densities may cause the radio synchrotron ages in the cavity backflow to be smaller than observed.  \\vskip.1in\\noindent$\\bullet$ If gas inside the radio lobe has a  small viscosity, as we assume, the KHI can be damped without strong magnetic fields.  \\vskip.1in\\noindent$\\bullet$ Radio lobe limb-brightening is due largely to the increased lobe magnetic field in the boundary backflow region. The wind that emanates from the shock spot carries non-relativistic gas, magnetic field and cosmic ray electrons, all of which flow back in the anti-jet direction along the radio lobe boundaries. The radio synchrotron emissivity depends on the product of the cosmic ray and magnetic energy densities, $e_cu_B$. However, if $e_c$ dominates the pressure in the radio lobes, as we assume, $e_c$ must be nearly uniform across (and throughout) the radio lobes to balance the pressure in the local external (shocked) cluster gas. Consequently, radio limb-brightening is not due to spatial variations in the CR electron density but instead to the varying magnetic field which is larger near the lobe boundaries.  Assuming a random sub-grid field morphology, the field anywhere in the lobe has adiabatically expanded from the value it had when it left the shock spot. But the field is frozen into the non-relativistic gas as it expands from the shock spot to a lower density in the observed lobe following $B \\propto \\rho^{2/3}$. The current gas density in the lobe that arises in the shock spot wind $\\rho_{lobe}(t)$ is approximately uniform inside the lobe. By assumption, the field in the shock spot remains constant with time, but the shock spot gas density $\\rho_{ss}$ decreases dramatically as the lobe evolves into the cluster gas. Because of this, the field in the lobe varies as $(\\rho_{ss}(t))^{-2/3}$ which increases with time during the lobe evolution and decreases with the age of the CRs. Therefore, the lobe field that emerged from the shock spot most recently (located in the boundary backflow) is larger than fields that entered the lobe at previous times. This enhances the field and synchrotron emissivity in the boundary backflow and explains the observed radio lobe limb-brightening. Evidently a similar argument also explains large-scale limb-enhanced magnetic fields and limb-brightening in lobes computed using the induction equation.  \\vskip.1in\\noindent$\\bullet$ It is shown how the analytic expression for the evolution of radio synchrotron electrons due to uniform expansion and radiation losses can be adapted to estimate the electron energy distribution $n(\\gamma,{\\bf r},t)$ and specific radio synchrotron emissivity $\\epsilon_{\\nu}({\\bf r},t)$ at every grid zone and time inside the radio cavity. These local functions are quite well determined from only three computed hydrodynamic arrays: gas density $\\rho$, magnetic field $B$ and the advected time since local synchrotron-emitting electrons left the post-shock region.  \\vskip.1in\\noindent$\\bullet$ For the first time the local energy spectrum $n(\\gamma)$ and the emissivity $\\epsilon_{\\nu}$ of radio-synchrotron electrons can be quantitatively computed throughout the radio cavity only from observations of the bright spot (when corrected to shock spot values) and simple assumptions about shock spot evolution over time.  \\vskip.1in\\noindent$\\bullet$ Parameters describing the double power law $n(\\gamma)$ observed in the Cygnus A bright spots need to be corrected to conditions in the nearby shock spot region. The shock spot is the source of electrons in the bright spot and the wind that inflates the radio cavity. This correction can be made from gasdynamical calculations.  \\vskip.1in\\noindent$\\bullet$ When this correction is made, the total energy density of relativistic synchrotron electrons in the current shock spot $e_{ce,ss}$ matches the total relativistic energy density $e_{c,ss}$ required to inflate the Cygnus A cavity to its current volume. This indicates that the relativistic fluid in radio cavities consists of electron pairs rather than a relativistic electron-proton plasma. A pure relativistic pair plasma (possibly with some entrained non-relativistic gas acquired during jet-gas interactions with the thermal filament) is physically attractive since it may be consistent with a purely electromagnetic origin of FRII jets (e.g. Blandford 2008).  \\vskip.1in\\noindent$\\bullet$ The computed radio synchrotron emission from Cygnus A cavities having purely toroidal magnetic fields is a very poor fit to VLA observations. Boundary backflows decelerate due to small negative radial pressure gradients inside the lobes. During deceleration, the toroidal field grows exponentially, causing the synchrotron emissivity to increase unrealistically along the backflow. Toroidal fields rapidly grow to exceed the fields observed in the Cygnus A radio cavity. Evidently the radial component of the field $B_r$, which is also transverse to the decelerating boundary backflow, would also increase in the same undesirable way along the backflow. This unfavorable outcome indicates that the usual MHD codes that solve the induction equation for the magnetic field on grid scales or larger are unlikely to provide realistic solutions inside radio cavities.  \\vskip.1in\\noindent$\\bullet$ Consequently, we consider a field that is random on subgrid scales and which advects as $B_{ran} \\propto \\rho^{2/3}$. Such a field may result from small scale turbulence due to plasma instabilities. When saturated, the magnetic energy density $u_B = B_{ran}^2/4\\pi$ may approach a low value comparable to the non-relativistic energy density $e$ which is generally less than the relativistic energy density $e_c$. Random subgrid activity would also be consistent with the viscous transport of non-relativistic gas, a likely source of viscous KH damping. Using a simple model for $B_{ran}$, we find that the radio synchrotron emissivity inside the Cygnus A cavity can be fit quite well. $B_{ran}$ is almost constant along the boundary backflow and has a value $\\sim 20\\mu$G that is almost identical to values observed in the radio cavity. The radio emissivity predicted by the random field model slowly decreases along the backflow, similar to VLA observations, with emissivities that quantitatively match these same observations.  \\vskip.1in\\noindent$\\bullet$ If the jet and shock spot parameters remain fixed, the separation between the shock spot and bright spot increases with time, becoming about three times larger when Cygnus A is at twice its current age. This can be explained by the reduced cluster gas density at larger distances from the cluster core, allowing the shock spot wind to evacuate a larger local volume.  \\vskip.1in"
    },
    "1402/1402.6969_arXiv.txt": {
        "abstract": "{Detailed spectral classifications are presented for 352 O--B0 stars in the VLT-FLAMES Tarantula Survey ESO Large Programme, of which 213 O-type are judged of sufficiently high quality for further morphological analysis. Among them, six subcategories of special interest are distinguished. (1) Several new examples of the earliest spectral types O2--O3 have been found, while a previously known example has been determined to belong to the nitrogen-rich ON2 class.  (2) A group of extremely rapidly rotating main-sequence objects has been isolated, including the largest \\vsini\\ values known, the spatial and radial-velocity distributions of which suggest ejection from the two principal ionizing clusters NGC~2070 and NGC~2060. (3) Several new examples of the evolved, rapidly rotating Onfp class show similar evidence, although at least some of them are spectroscopic binaries. (4) No fewer than 48 members of the Vz category, hypothesized to be on or near the zero-age main sequence, are found in this sample; in contrast to the rapid rotators, they are strongly concentrated to the ionizing clusters and a newly recognized region of current and recent star formation to the north, supporting their interpretation as very young objects, as do their relatively faint absolute magnitudes.  (5) A surprisingly large fraction of the main-sequence spectra belong to the recently recognized V((fc)) class, with \\cc\\ emission lines of similar strength to the usual \\nc\\ in V((f)) spectra, although a comparable number of the latter are also present, as well as six objects with very high-quality data but no trace of either emission feature, presenting new challenges to physical interpretations. (6) Two mid-O~Vz and three late-O giant/supergiant spectra with morphologically enhanced nitrogen lines have been detected.  Absolute visual magnitudes have been derived for each star with individual extinction laws, and composite Hertzsprung-Russell Diagrams provide evidence of the multiple generations present in this field.  Spectroscopic binaries, resolved visual multiples, and possible associations with X-ray sources are noted.  Astrophysical and dynamical analyses of this unique dataset underway will provide new insights into the evolution of massive stars and starburst clusters.} ",
        "introduction": "\\label{sect: intro}  30~Doradus in the Large Magellanic Cloud (LMC) comprises the most massive starburst cluster and giant \\hb\\ region in the Local Group, and it contains the most massive stars known (Crowther et al. 2010; Bestenlehner et al. 2011).  It is the paradigm for understanding early massive stellar and cluster evolution, and for interpreting more distant starbursts.  Thus it is appropriate that 30~Dor is the subject of the unprecedented spectroscopic dataset obtained by the VLT-FLAMES Tarantula Survey (VFTS; ESO Large Programme 182.D-0222; Evans et al. 2011), which contains high-resolution ($R \\sim 10^4$) observations of about 800 OB stars in this field.  Of these, 352 from the Medusa-Giraffe multi-object spectrograph configuration have been determined to be of spectral type O (or B0 in 10 cases included here) and are the subject of this morphological investigation.  The prior state of the art in digital spectral classification of the O stars is represented by the Galactic O-Star Spectroscopic Survey (GOSSS; Ma\\'{i}z Apell\\'aniz et al. 2011; Sota et al. 2011), in which the system and procedures have been further developed and refined. An expanded list of standard stars and an extensive new classification atlas from those high-S/N data are presented there; many of the GOSSS developments have been applied in the present work.  However, the resolution of the VFTS data is substantially higher, which has required additional developments and another new atlas of primarily Galactic spectra, from the IACOB survey (Sim\\'on-D\\'{\\i}az et al. 2011) in the northern hemisphere and the ESO archives in the southern, as discussed in detail and presented by Sana et al. (2014a, in prep.).  Most of that discussion will not be repeated here, but it was determined that $R \\sim 4000$ is the practical upper limit for reliable visual spectral classification of O-type spectra, because of the effects of line-profile differences and numerous weak lines at higher resolving powers, so  that value has been adopted for the new atlas and for this work.  Section~2 provides some further details about the data and classification procedures used here, including the specification of the high-quality subset of 213 objects isolated for further analysis, based on not only data characteristics but also multiplicity considerations.  Experience has shown that substantial advances in data quality and/or sample size generally produce both foreseeable and unexpected new scientific results; the present study is no exception, and Section~3 describes six spectroscopic categories of special interest encountered among this subsample, some previously recognized (with new members added here) but others newly distinguished. Section~4 references the essential, tailored extinction corrections required for astrophysical analysis in regions of this kind and presents empirical and theoretical  Hertzsprung-Russell Diagrams (HRDs) for the high-quality subsample.  Possible associations with {\\it Chandra} X-ray sources (Townsley et al. 2006) are noted in Section~5.  Section~6 provides a summary, conclusions, and outlook for further analyses of these data by the VFTS Team. ",
        "conclusions": "This work provides basic spectroscopic descriptive information to complement and support a wide array of analyses undertaken by the various specialists of the diverse VFTS Team.  Specifically, detailed spectral classifications have been presented for 352 O and B0 stars, many of which have been found to belong to several categories of special interest.  These include the earliest (O2--O3) spectral types, of which nine new members have been found and for an equal number of which refined types are given.  A salient result is the detection of a group of 18 very rapidly rotating O main-sequence stars, including the fastest known to date with \\vsini\\ $\\sim600$~km~s$^{-1}$, for which radial-velocity and spatial distributions support a hypothesis of ejection from the 30~Dor clusters. Other noteworthy objects include new members of the also rapidly rotating but evolved Onfp class, and a surprisingly large number (48) of Vz objects, likely related to the lower metallicity and weaker winds in the LMC.  They have been hypothesized to be in a very early evolutionary state, which is substantiated by their locations both on the sky and in the HRDs, although further interesting aspects of their interpretation have emerged from parallel quantitative analysis as cited below.  Further intricacies and diversity of O-type main-sequence spectral morphology have been uncovered, including 19 members of the O V((fc)) class with C~III emission features equal to or stronger than the usual N~III, as well as six with very high S/N but no trace of either emission species; this complexity will no doubt yield significant physical diagnostics when modeled and understood.  Finally, several objects with morphological anomalies in their CNO spectra have been found, which are related to massive stellar evolution and mixing or transfer of processed material.  Basic information about VFTS has been provided by Evans et al. 2011 (Paper~I), and further global analyses of these data from different perspectives have appeared or are in progress.  Multiplicity among the O-type sample based on radial velocities has been investigated by Sana et al. (2013; Paper~VIII) with the result that the spectroscopic binary fraction is at least 50\\%.  The interstellar content of the spectra has been extracted by van Loon et al. (2013; Paper~IX), including the detection of high-velocity absorption-line components in \\naa; while the nebular emission-line data have been presented by Torres-Flores et al. (2013), revealing several new kinematic structures in the nebula.  A census of the hot stars and their feedback, incorporating the present spectral classifications, has been compiled and discussed by Doran et al. (2013; Paper~XI).  The rotational velocities of the single O stars have been derived by Ram\\'{\\i}rez-Agudelo et al. (2013; Paper~XII).  The physical properties of the Vz class have been investigated by Sab\\'{\\i}n-Sanjuli\\'an et al. (2014; Paper~XIII), providing important insights into its atmospheric and evolutionary significance. Detailed atmospheric and wind analyses of 62 VFTS O, Of/WN, and WN stars are presented by Bestenlehner et al. (2014; Paper~XVII).   Ultimately, exhaustive quantitative stellar spectral analysis of the full VFTS sample (complemented by {\\it HST} spectroscopy in R136, PI P.~Crowther; proper motions, PI. D.~Lennon; photometry, Sabbi et al. 2013), as well as studies of extinction (Ma\\'{\\i}z Apell\\'aniz et al. 2014; Paper~XVI) and nebular lines, will support theoretical analyses of massive stellar evolution and the nebular environment, by means of which this unprecedented dataset for the unique 30~Doradus region will advance understanding of both massive stars and their starburst habitats."
    },
    "1402/1402.4477_arXiv.txt": {
        "abstract": "Wolf-Rayet  stars are  known  to eject winds.  Thus, when  a Wolf-Rayet  star explodes  as a  supernova, a  fast, $>30,000$  km/s, shock  is expected to be driven  through a wind. We study the  signal expected from a fast  supernova shock  propagating through  an optically thick  wind, and  find that  the electrons behind  the shock driven  into the  wind are cooled  efficiently, by inverse  Compton  over soft  photons  that  were  deposited by  the  radiation mediated shock that crossed the  star. Therefore, the bolometric luminosity is comparable to the  kinetic energy flux through the shock,  and the spectrum is found to be a power-law, which slope  and frequency range depend on the number flux of soft photons available for cooling. Wolf-Rayet supernovae that explode through  a thick  wind have  a high  flux of  soft photons,  producing a  flat spectrum, $\\nu F_{\\nu}=Const$,  in the X-ray range  $0.1\\lesssim T\\lesssim 50$ keV. As the shock expands into an optically thin wind, the soft photons are no longer able to  cool the shock that  plows through the wind,  and the bulk of the emission takes  the  form of  a  standard  core-collapse  supernova (without  a  wind). However, a  small fraction of the  soft photons is upscattered  by the shocked wind     and    produces     a    transient     unique    X-ray     signature. ",
        "introduction": "The scenario of a supernova (SN) exploding through a thick wind has been the focus of many recent papers \\citep[e.g.][]{Ofek+2010,Balberg_Loab2011,Chevalier_Irwin2011,Murase+2011,Katz+2011,Moriya_Tominaga2011,Chevalier_Irwin2012,Svirski+2012,Balberg+2012}. The events discussed in these papers are mostly classified as type IIn SNe, and although they only account for a few percents of the observed SNe (e.g. \\citealt{Li+2011}), they provide a unique probe for the progenitor's pre-explosion composition and mass-loss history. In some cases, such information may even reveal an unexpected new SN progenitor type (e.g., a luminous blue variable, \\citealt{Kiewe+2012,Mauerhan+2013}) or a possible causal connection between a pre-explosion mass-loss burst and the SN explosion \\citep{Ofek+2013}.   The study  of SNe  that interact  with a thick  wind has  focused thus  far on interaction  timescales of  days or  longer, currently  the most  probable for detection. However,  some Wolf-Rayet (WR)  SNe are  naturally expected  to interact with a thick wind, and this interaction timescale should be much shorter than a   day.    Here   we   study    this   common   yet    unexplored   scenario.  If, due to a massive mass loss, a SN progenitor is surrounded by a wind of an optical depth $>c/v$, where $c$ is the speed of light and $v$ is the SN shock velocity, then the shock does not break out near the stellar edge: After crossing the star's envelope, the radiation dominated shock continues into the wind, while a reverse shock is driven into the SN ejecta. The interaction of the expanding ejecta with the thick wind leads to a very luminous SN, that breaks out once the optical depth of the wind ahead of the shock decreases to $c/v$. Following the breakout, the radiation mediated shock is replaced by a narrower collisionless shock and a layer of hot shocked electrons, initially with a temperature $T_h\\gtrsim 60$ keV (hereafter, a temperature $T$ denotes an energy $k_BT$), forms behind the shock \\citep{Katz+2011,Murase+2011}. The interplay of this layer with the surrounding colder medium, mediated by the diffusing photons, regulates the luminosity and determines the observed spectrum.  In \\cite{Svirski+2012} we discussed the observed signal from events with breakout times of a few days after the SN explosion or longer. For standard SNe parameters these events map to a breakout shock velocity $v_{bo,9}\\le1$, in units of $10^9 {\\rm~cm~s^{-1}}$. The analysis in \\cite{Svirski+2012} is limited to scenarios in which the radiation emitted by the cold medium that surrounds the hot layer is in thermal equilibrium with the emitting electrons. Svirski et al. (in preparation) discuss non-equilibrium scenarios.  Here we turn  our focus to a yet  unexplored regime of interacting  SNe, that of extremely high  shock velocities  found in  compact progenitor  explosions. WR progenitor SNe are  expected to develop $v_{bo,9}\\gtrsim  4$, corresponding to breakout times of minutes rather than days (see Equation \\ref{EQ v_bo} below). In addition,  WR stars  go through a massive  mass loss and  thus  some are likely to explode through a  thick wind. We find  that the emission from  such events is X-ray  dominated, and  therefore their  detection is  challenging both  due to their short duration and due to the  narrow field of view that current sensitive X-ray detectors  provide. However,  the serendipitous  X-ray detection  of SN  2008D \\citep{Soderberg+2008},  a  SN of  an  extreme  shock velocity  that  possibly exploded through a  thick wind, encourages us to study  the parameter space of such                                                                   events.  If the progenitor of a core-collapse SN  is \\emph{not} surrounded by a wind, then the early SN luminosity is powered by the internal energy that the SN shock deposits in the progenitor's envelope. If, in contrast, a thick wind \\emph{does} surround the progenitor, then the interaction of the expanding ejecta with the wind dominates the event luminosity (at least as long as the unshocked wind is optically thick). In this case, the radiation deposited in the envelope by the shock, rather than escaping freely to the observer, diffuses further through the interaction layer and affects its cooling. Although the energy flux of this radiation may be considerably lower than the flux powered by the interaction, we find that it has a profound influence on the observed luminosity and spectrum of a WR exploding through a thick wind. While a failure to account for this radiation results in a faint monochromatic signal of $T \\gtrsim 60$ keV, its proper inclusion invokes an efficient cooling of the shocked plasma by inverse Compton interactions and implies a much higher luminosity -- comparable to the shock kinetic energy flux, and a rather flat, $\\nu F_{\\nu}= Const$, observed spectrum that spans over a large frequency range that ends at $T<60$ keV.  Our analysis applies to SNe that explode through a thick wind and reach a breakout velocity $v_{bo,9}\\gtrsim 3$, including highly energetic explosions in blue supergiant and red supergiant progenitors, in cases that they are surrounded by a thick wind. While motivated by SNe that explode through a thick stellar wind, the solution that we provide is general and may apply to other physical settings that involve a fast collisionless shock in a thick medium.  In Section \\ref{sec:analysis} we develop a general solution for the spectrum of a fast collisionless shock propagating through an optically thick medium, in the presence of an external soft radiation source. In Section \\ref{sec:impl} we review the implications of our analysis for fast shock SNe exploding through a thick wind. In Section \\ref{sec:post} we discuss the evolution of such systems as the wind ahead of the shock becomes optically thin. We summarize in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We derive the spectrum and luminosity observed from a fast, $> 30,000$ km/s SN collisionless shock propagating through an optically thick wind, and find that the injection of soft photons into the interaction region plays the main role in shaping the spectrum and light curve. While previous works focused at lower shock velocities, $\\lesssim 10,000$ km/s, and provided only rough predictions regarding the observed spectrum \\citep[e.g.][]{Chevalier_Irwin2012,Svirski+2012}, here we elaborate on these works and provide, for the first time, a detailed prediction for the observed spectrum. Fast SN shocks propagating through a thick WR wind develop a flat spectrum with an initial frequency range $0.1\\lesssim T\\lesssim 50$ keV that becomes wider with time. If a SN explosion of a red or a blue supergiant drives a fast shock and the shock traverses a thick wind, the spectrum depends on the ratio $\\frac{R_*}{R_{bo}}$, which value cannot be constrained a-priori.  In detail, we solve for the observed signal as a function of $f_{inj}$, $T_{inj}$, $t_{bo}$ and $v_{bo}$. When the shock is fast cooling, $f_{inj}$ and $T_{inj}$ determine the spectrum at $\\tau\\gtrsim 1$, and assuming  a standard wind density profile ($\\rho\\propto r^{-2}$), $t_{bo}$ and $v_{bo}$ determine the luminosity evolution. First, we find the threshold conditions for a fast cooling of the shock, Equations \\ref{eq:condition} and \\ref{eq:condition2}. When these conditions are satisfied, solving Equations \\ref{eq:energy_balance}, \\ref{eq:slopes2} and \\ref{eq:Tmin} provides the spectral slopes and ranges. Generally, at lower energies, the spectrum is a power-law $\\nu F_{\\nu}\\propto \\nu^{\\alpha}$ of $0\\lesssim\\alpha<1$, and it typically has a break above which $\\alpha=-f_{inj}$. The spectrum spans from $T_{min}\\gtrsim T_{inj}$ up to $m_ec^2/\\tau^2\\lesssim T_{max}\\lesssim T_h$, where the exact $T_{min}$, $T_{max}$ and the break temperature are found by solving the equations as described above. An injected photon number flux mildly above the fast cooling threshold yields a hard spectrum, while a higher flux yields a softer spectrum. Over a wide range of injection parameters, the spectrum that evolves is approximately flat, $\\nu F_{\\nu}=Const$. The condition for a flat spectrum is given by Equation \\ref{eq:flat2}. Section \\ref{sec:picture} provides a brief description of the physical processes that take part in shaping the spectrum. When propagating through a \\emph{thin} wind, the interaction signal is sub-dominant but it provides a unique signature through its X-ray luminosity decay pattern, $L_X(t)\\propto t^{-n}$, where $n$ is the number of collisions with the hot layer that brings an Optical photon up to the X-ray detection window.  \\begin{figure*} \\centering \\epsscale{1} \\plotone{conclusion6.eps} \\caption{A schematic evolution of the X-ray luminosity (main figure) and the spectrum (inset) from a SN shock breakout through a thick WR wind, sketched for $t_{bo}=100$ s, $\\tau_{bo}=5$, $f_{inj}=0.1$ and an initial $T_{inj}=0.1$ keV. (A) is the breakout pulse, with a non-thermal spectrum peaking at a few keV. (B) is the phase of IC fast cooling over injected photons, matching $\\tau\\gtrsim 1$, with $L\\propto t^{-0.4}$ and a nearly flat spectrum that gradually widens from B1 to B2. (C) is the luminosity drop by $\\sim 1/f_{inj}$ that follows the transition to a slow cooling at $\\tau\\lesssim 1$, and (D) reflects $\\tau\\ll 1$, with $L_X\\propto t^{-3}$ and a negative spectral slope that steepens with time.} \\label{fig:evolution} \\end{figure*}  After deriving our model we apply it to infer the observed spectrum and light curve from a SN shock breakout through a thick WR wind. Figure \\ref{fig:evolution} depicts the different evolution stages of the X-ray luminosity and the spectrum for typical WR parameters, $t_{bo}=100$ s, $v_{bo,9}=6$, $\\tau_{bo}=5$, $f_{inj}=0.1$ and an initial $T_{inj}=0.1$ keV. The first radiation observed is the breakout pulse, with a rise time $\\sim 100$ s (A in Figure \\ref{fig:evolution}) and a non-thermal spectrum peaking at a few keV (A in the inset). The collisionless shock that forms after the breakout is cooled efficiently and the bolometric luminosity decays slowly, $L\\propto t^{-0.4}$ (B in Figure \\ref{fig:evolution}). The observed spectrum is nearly flat, with an initial relatively narrow range in the X-rays (B1 in the inset) that gradually widens beyond X-rays as $\\tau\\rightarrow 1$ (B2 in the inset). At $\\tau<1$ the shock cooling becomes slow, and the X-ray luminosity drops by a factor $\\sim 1/f_{inj}$ within a dynamical timescale (C in Figure \\ref{fig:evolution}). Later, when $\\tau\\ll 1$, the X-ray luminosity is dominated by $\\lesssim$ keV photons and decays rapidly, $L_X\\propto t^{-3}$ (D in Figure \\ref{fig:evolution}).  The solution we provide for the thick wind phase is general and may apply to other physical settings involving an irradiated fast collisionless shock that propagates through an optically thick medium.  G.S. and E.N. were partially supported by an ISF grant (1277/13), an ERC starting grant (GRB-SN 279369), and the I-CORE Program of the Planning and Budgeting Committee and The Israel Science Foundation (1829/12).  \\appendix"
    },
    "1402/1402.3257_arXiv.txt": {
        "abstract": "{The understanding of Galaxy evolution can be facilitated by the use of population synthesis models, which allows us to test hypotheses on the star formation history, star evolution, and chemical and dynamical evolution of the Galaxy. } {The new version of the Besan\\c{c}on Galaxy model (hereafter BGM) aims to provide a more flexible and powerful tool to investigate the initial mass function (IMF) and star formation rate (SFR) of the Galactic disc.} {We present a new strategy for the generation of thin disc stars, which assumes the IMF, SFR and evolutionary tracks as free parameters. We have updated most of the ingredients for the star count production and, for the first time, binary stars are generated in a consistent way. The local dynamical self-consistency is maintained in this new scheme. We then compare simulations from the new model with Tycho-2 data and the local luminosity function, as a first test to verify and constrain the new ingredients. The effects of changing thirteen different ingredients of the model are systematically studied.} {For the first time, a full sky comparison is performed between BGM and data. This strategy allows us to constrain the IMF slope at high masses, which is found to be close to 3.0 and excludes a shallower slope such as Salpeter's one. The SFR is found decreasing whatever IMF is assumed. The model is compatible with a local dark matter density of 0.011 $M_\\odot$pc$^{-3}$ implying that there is no compelling evidence for the significant amount of dark matter in the disc. While the model is fitted to Tycho-2 data, which is a magnitude limited sample with V$<$11, we check that it is still consistent with fainter stars.} {The new model constitutes a new basis for further comparisons with large scale surveys and is being prepared to become a powerful tool for the analysis of the Gaia mission data.}% ",
        "introduction": "The understanding of the origin and evolution of the Milky Way is one of the primary goals of the Gaia mission (ESA, launched December 2013). The use of its data to test different hypotheses and scenarios of galaxy formation and evolution requires the availability of an adaptable Galaxy model to provide simulated data for comparison. Kinematic and photometric data with the ages and metallicities of the stars and their statistics allows us to characterise galaxy's populations and, from that, the overall Galactic gravitational potential.  One of the most promising procedures to reach such a goal is to optimize the present population synthesis models by fitting robust statistical techniques, the large and small scale structure, and kinematics parameters that best reproduce Gaia data. The work presented in this paper is focused on the optimization of the structure parameters of the Milky Way Galactic disc in the BGM by comparing the simulations to real data to study the process of Galaxy evolution.  The development of Galaxy models started in the early 80's when several needs for star count predictions for the sake of preparing observations, in particular HST mission. For example, \\cite{Bahcall} built a model based on simple assumptions on density laws and luminosity functions, while \\cite{Cohen} developed a model in the infrared using a more complex set of luminosity functions but leading to a very large number of free parameters. From the beginning, the aim of the BGM was not only to be able to simulate reasonable star counts but further to test scenarios of Galactic evolution from assumptions on the rate of star formation, initial mass function, and stellar evolution. This was the spirit of the \\cite{Tinsley1972A&A....20..383T} formalism of decomposition of the star-creation process in these two functions of mass and time, which are assumed independent. \\cite{Creze1979} and \\cite{CrezeRobin1983} conceived the starting point of a model to be able to generate samples of stars directly comparable with observations. The idea was that simulations can be efficiently compared with data because all observational bias can easily be taken into account. Then statistical tests and fitting processes can be applied to constrain model parameters and hypothesis, such as the IMF, the density laws, etc. The initial BGM from \\cite{ar1986} was improved for taking into account dynamical constraints (see \\cite{bienayme}). Further developments concerned the stellar and Galactic evolution constraints from the Hipparcos input catalogue, the high latitude star counts (\\cite{misha1997-2}), constraints on the thick disc population \\citep{Robin1996,Reyle2001}, the halo (\\cite{Robin2000}), and the bulge \\citep{Picaud2004}. This leads to a version, which was released to the community on the web \\footnote{http://model.obs-besancon.fr} and described in \\cite{ar2003}. Since this work, we continued to develop it, partly for giving improved fit to available data and partly to constrain Galactic structure and scenarios of galaxy formation, especially in the Galactic plane, thanks to the 3D extinction model developed \\citep{marshall3D2006}, the study of the Galactic warp and flare \\citep{Derriere,Reyle2009}, and, lastly, the bar and bulge regions \\citep{ar2012}.  In recent years, other population synthesis models have been developed, such as those by \\cite{NG}, \\cite{Vallenari}, \\cite{Girardi2005}, and that known as tri-dimensional model of the galaxy (TRILEGAL). These models are also based on stellar evolutionary tracks and assumptions on SFH and IMF. However, none of them have yet used dynamical constraints, as BGM has. The model TRILEGAL benefits for a wide coverage of all the photometric systems used in observations. It uses slightly different model parameters; the atmosphere model used are the one from Padova, while we use semi-empirical grids from BaSeL2.2 \\citep{Lejeune1997,Lejeune1998} or 3.1 \\citep{Westera2002} complemented at low temperatures by NextGen models in BGM. The colour predictions are then significantly different in the low mass region. % The TRILEGAL model uses essentially the same kind of parametrization for the thin disc with a scale height changing with age but with an exponential or secant squared while BGM uses Einasto ellipsoids. A simple diffuse extinction model from an exponential or sech$^{2}$ function is used for modelling the interstellar absorption, and the warp and flare in the outer disc are not yet taken into account in TRILEGAL. Bulge parameters, as proposed by \\cite{vanhollobeke} are sensitively different than in BGM \\citep{ar2012}. \\cite{NG} and \\cite{Vallenari} models have only been compared to data in a limited number of directions of observations, and their parameters are less well established and constrained.  \\cite{JJ2010} proposed a new scheme, using local dynamical constraints to create a local model of the Galaxy. This model only accounts for main sequence stars and have not been compared with large scale surveys, but it produced interesting constraints on the local star formation history and on the IMF. \\cite{Gao} attempted to combine the \\cite{JJ2010} approach with TRILEGAL code to fit some parameters and improved the fit to SDSS data towards high latitudes. However, they conclude that TRILEGAL has not enough flexibility to obtain a satisfactory fit to these data and show that BGM (2003 version) does not compare well with SDSS data. However, this was a version of the BGM that was not calibrated for the ugriz photometric system yet.   Galaxia (\\cite{Sharma}) is another type of approach, which makes use of the BGM (the 2003 version) in order to create large simulations to be compared with large scale surveys. It can take any Galaxy model as an input but is not a new Galaxy model by itself. It is also a very useful tool to translate N-body simulations into an observable catalogue.  For a long time, we have identified a systematic disagreement between BGM and real data concerning bright stars. A comparison to the Hipparcos catalogue showed that the A-F dwarfs are over-represented by 20-30\\% while the quality of star count predictions are at the level of a few percent at fainter magnitudes (\\cite{ar2003,IbataI}). It appears that faint star counts (typically  V$\\ge$ 15) are dominated by thick disc and halo populations at or slightly below the turnoff, while the contribution of the thin disc is dominated by K and M dwarfs, a region of the HR diagram where the history of star formation is not important. With these stars being unevolved, this occurs because their density depends on the integral of the SFR over the age of the disc and not on the detailed history of the SFR.  Since the study of \\cite{misha1997-1} on the Hipparcos input catalogue, which showed that the SFR should be nearly constant or not varying by more than a factor of 3 over the age of the thin disc, we assumed a constant star formation for the thin disc in BGM. The A-F dwarfs, which constitute a large contribution to bright data sets, as Hipparcos and Tycho-2, are indeed very sensitive to the history of the SFR, and we concluded that it was necessary to specifically study this aspect using these dwarfs in the solar neighbourhood to get rid of the discrepancy of the BGM. In doing so, the Tycho-2 catalogue is the best homogeneous data set available for the whole sky and complete to about magnitude $V \\sim 11$. This is the specific subject of the present study.  In the era of Gaia, a model able to analyse stellar population distributions to constrain Galaxy evolution would be very welcome. The BGM is in this respect a very useful tool and should be able to ``measure'' the star formation history in different regions of the Milky Way. However, the old scheme of BGM, based on the evolution scheme from \\cite{misha1997-1} does not have enough flexibility to change the IMF and SFR in the thin disc. Moreover, it does not allow us to generate binary stars easily. This is why we have undertaken the present work to modify the BGM code to be able to change these parameters easily to use different ones in different regions of the Milky Way if necessary and to account for stellar binarity at the same time. To test this scheme and to constrain the star formation history, the new code presented here has been compared and tested using the Tycho-2 catalogue and the local luminosity function (LF). We consider this as a first step for furnishing a reliable model for interpretation of future large scale surveys, such as Gaia, but also for RAVE, APOGEE, Gaia-ESO, and other future data.   In Section \\ref{bgmCode}, the new approach and the overall structure of the code are presented % . In Section \\ref{tools}, we present the observables and tools used in our analysis. In Section \\ref{sec:ingred}, we detail the model ingredients, which have been investigated. In Section \\ref{res}, we show the results obtained from the multiple comparisons of our model with the Tycho-2 data and the LF; additionally it is our first check of how the new model reproduces deep star counts. ",
        "conclusions": "\\label{summary}   We have designed, implemented, and tested a new version of the Besan\\c{c}on Galaxy Model. The new code implements an important change in the star production philosophy of the thin disc population, where the IMF, SFR, and stellar evolutionary tracks are treated as free parameters. The new code can be tested by comparing simulations by assuming different IMF and SFR, which could also be different in different parts of the Galaxy. The code also produces binary systems that take into account the observational spatial resolution.  We have updated and analysed several of the model's inputs (see Table \\ref{tableParam}) using two observational sets, the Tycho-2 data, and the local luminosity function.  The extinction model is one of the most  crucial inputs at low latitudes. Significant differences in the Galactic plane are observed when using \\cite{drimmel2001} or \\cite{marshall3D2006} extinction models. Our results indicate that \\cite{drimmel2001} overestimates the extinction at large distances and \\cite{marshall3D2006} underestimates the extinction at very short distances. In absence of a better alternative, we have selected the \\cite{marshall3D2006} whenever it is available and the \\cite{drimmel2001} for the rest of the sky.  One of the most crucial parts of our investigation was to look for the best IMF and SFR, which are able to reproduce the whole sky Tycho-2 data. When a constant SFR  is assumed, no matter which IMF is used, we are not able to reproduce the Tycho-2 star counts. The model always shows a significant excess of blue stars. This result strongly favours a decreasing SFR. The new decreasing SFR adopted here does not alter the overall good agreement obtained with the previous model (constant SFR) at fainter magnitudes. This is due to the fact that fainter stars (V$>$14) of the thin disc are on the lower main sequence where the stellar density is sensitive to the integral of the SFR over the disc age and not to the detailed evolution in time.   When looking for the best fit IMF, we notice the following:  \\begin{itemize} \\item A slope $\\alpha$=2.3 at high masses, which has been used by \\cite{Vallenari2006} and \\cite{Kroupa2008} produces a strong excess of blue stars. As already mentioned by \\cite{misha1997-1}, a Salpeter IMF is certainly inappropriate for our Galaxy. A steeper slope in this mass range is recommended. \\item  The \\cite{JJ2010} IMF used with their best fit decreasing SFR, and our default model reproduces the Tycho-2 $(B-V)_T$ distribution at intermediate and high latitudes well, but it performs a bit worse within the plane. \\end{itemize}  We conclude that Kroupa-Haywood v6 and Haywood-Robin IMF combined with the decreasing SFR of \\cite{AumerBinney2009} provide the best fit to Tycho data.  All the thin disc ingredients cannot be constrained using nearby photometric data alone. A better solution would require a higher amount of observational constraints (in particular deeper star counts) and a robust statistical model fitting algorithm (as for instance Markov Chain Monte Carlo method). Nevertheless, we have learned that Tycho data is able to impose strong observational constraints to the SFR and the IMF in the bright star/intermediate mass domain in the solar vicinity.  A complementary investigation of the thick disc is underway using SDSS and 2MASS surveys to characterise the thick disc properties and constrain the scenario of its formation. We believe that using this new scheme and a better knowledge of the thin disc itself will help to perform this analysis, getting rid of the uncertainties on the structural parameters due to the degeneracies of parameters and to the thin disc contamination. Scenarios, such as the decomposition of populations of mono-abundances from \\cite{Bovy}, the question of the role of migration, in situ formation, or formation by mergers will then be possibly tested.  The next step will be to test the stellar kinematics and metallicity link to constrain the chemo-dynamical evolution of the Galaxy using RAVE data. The availability of even deeper spectroscopic probes, like the Gaia-ESO survey or APOGEE, will also allow us to constrain the SFR on a larger scale and to specifically study the link between the local thin and thick discs and more problematic populations of the bar and the bulge in the inner Galaxy. We expect that the new tool presented here will be very efficient for the analysis and understanding of future large scale surveys, such as Gaia. Work is in progress to make the new default models available online."
    },
    "1402/1402.6534_arXiv.txt": {
        "abstract": "The \\ik mission has discovered over 2500 exoplanet candidates in the first two years of spacecraft data, with approximately 40\\% of them in candidate multi-planet systems.  The high rate of multiplicity combined with the low rate of identified false-positives indicates that the multiplanet systems contain very few false-positive signals due to other systems not gravitationally bound to the target star (Lissauer, J.~J., et al., 2012, ApJ 750, 131).   False positives in the multi-planet systems are identified and removed, leaving behind a residual population of candidate multi-planet transiting systems expected to have a false-positive rate less than 1\\%.  We present a sample of 340 planetary systems that contain 851 planets that are validated to substantially better than the 99\\% confidence level; the vast majority of these have not been previously verified as planets.  We expect $\\sim$2 unidentified false-positives making our sample of planet very reliable.  We present fundamental planetary properties of our sample based on a comprehensive analysis of \\ik light curves and ground-based spectroscopy and high-resolution imaging.  Since we do not require spectroscopy or high-resolution imaging for validation, some of our derived parameters for a planetary system may be systematically incorrect due to dilution from light due to additional stars in the photometric aperture.  None the less, our result nearly doubles the number of verified exoplanets. ",
        "introduction": "Data from the first two years of \\ik spacecraft operations have identified 3670 target stars with periodic or transit-like signatures indicative of transiting planets or eclipsing binary stars.  Approximately 50\\% of these targets have signatures that can be attributed to false-positives (FPs), primarily eclipsing binaries (EBs) centered on the target star, a chance alignment of a distant EB within the photometric aperture, or flux bleeds into the photometric aperture.  The remaining 2530 systems are composed of primarily exoplanetary systems with an expected FP rate of approximately 10\\% due to photometric blends \\citep{mor11,fre13,san13}.  However,  a subset of 457 systems show more than one candidate transiting planet candidate (PC); we refer to these candidates as ``multis\". FPs should be nearly randomly distributed among \\ik targets, whereas if flat multi-planet systems are common, then many targets should have multiple transiting planets \\citep{lis11}.  The large number of multis observed thus implies a high reliability rate, as quantified by (\\citealt{lis12,lis13}; henceforth Paper I and Paper II, respectively).  The small number of FPs found in multis observed by \\ik \\citep[and \\S\\ref{vps} of this paper]{lat11} reinforces our confidence in the high reliability of the PCs remaining in multis (see Appendix C of Paper II for details). After making selections to minimize the odds of blend scenarios, we find 340 systems containing a total of 851 planets that can be validated to better than the 99 percentile, with 768 planets across 306 systems being newly validated.  Some of these systems have also been confirmed from radial velocity detection \\citep{mar13,gau12}, transit timing variations \\citep{for11,ste12} and planet validation techniques such as {\\it BLENDER} \\citep{tor11,fre12} and now through multiplicity boost (Paper I).  We increase the known number of exoplanets from 942\\footnote{Based on NASA Exoplanet Archive 2013/11/12} to 1710.  With excellent precision and high-duty cycle, \\ik observations of transiting exoplanet systems provide photometric data that can be used to measure fundamental transit properties such as transit duration and observed transit depth.  These properties are determined by the relative sizes of the planet, host star and additional flux sources (other stars) within the photometric aperture.  Since observations provide nearly full coverage of the planetary orbit, the resulting photometric phase curves enable a useful diagnostic for the identification of astrophysical FPs that can mimic an exoplanet transit signature.  In \\S \\ref{lca} we examine the light curves of the \\ik sample and describe the nature of the planetary systems and how they are identified.  As an imaging instrument, \\ik also provides time-series measurements of the centroid of the photometric signal.  When multiple sources are present in the photometric aperture, the photometric centroid can move in response to flux changes from any of the sources.  This property allows \\ik pixel level data to be used to search for scenarios where a planetary transit-like event is produced by a diluted background eclipsing binary star \\citep{bat10}.  Paper II presents a theoretical exploration of the expected and predicted FP rate for transiting multi-planet systems.    From the $\\sim$190 000 targets observed by \\ikt, there are roughly 2500 transit-like patterns of events in the KOI catalogue, split between PCs and EBs.  Thus, for a \\ik target chosen at random it is unlikely that a transit event will be present.  One of the most common classes of FP event is caused by a background eclipsing binary (BEB) in the aperture of the target star.  The source of a BEB can either be an additional star that is found within or close to the photometric aperture or a bright star within the \\ik Field of View (FOV) that introduces flux in the photometric aperture due to optical ghosting, such as mirror images, or electronic interference, such as CCD crosstalk \\citep{cou13}.   The occurrence rate of FPs is largely independent of the target stars and thus it is far more likely that a FP source will produce a single transit-like event as opposed to a photometric light curve containing transit-like signals from multi-planet sources.   It is important to distinguish between FPs produced by background stars from those caused by instrumental effects.   We identify the latter as period/phase (P/T0) collisions to indicate that the period and epoch solution for the candidate event are not due to a unique event.  Rather, the transit signal from one target is seen on another target as well.  P/T0 collisions account for most of the identified instrumental FPs.  A full description and catalogue of P/T0 collisions can be be found in \\citet{cou13}.  Figure 1 of Paper II shows the galactic distribution of \\ik targets, PCs and FPs.  The results indicate that only a small fraction of the remaining \\ik PCs are likely to be background eclipsing binary (BEB) FPs.  The process of identifying and cataloguing FPs from the KOI list continues to evolve.  Thus, the FP list from various iterations of the KOI catalogue \\citep{bor11,bat12,bat13} have different reliability rates.  This paper describes in detail the steps taken to develop a reliable and uniform classification scheme and its application to the multi-planet sample to classify KOIs as false-alarms (FAs), FPs or PCs.  FAs are transit event candidates with a signal-to-noise ratio (S/N) below a S/N threshold of 7.1 (see \\S \\ref{snr}) or a transit candidate mimicked by stellar variability or an instrumental artifact.  We find that the FP rate for multi-planet systems is low and consistent with the predictions from Appendix C of Paper II based on a statistical analysis of the FP rate found in the single planet population.  The predicted FP rate allows us to conclude that the PCs that pass our FP and morphological tests would misclassify only $\\sim$2 PCs, allowing us to claim that 768 PCs are bona fide planets with a confidence level greater than 99\\%.  As an example, consider only the non-transiting, transiting single planet and transiting double planet systems.  Here we ignore the case of FP+FP, where two FPs are associated with the same target and estimate the number of FP for the double planet systems in the spirit of Paper II.  There are approximately 140 000 stars that contribute to the transiting planet population (\\S 4 of Paper II).  After removing FAs and T/P0 collisions there are 2182 systems that show one transiting body and 284 systems that show two transiting bodies.  From the single-planet systems, 662 systems were identified as FPs, which provides an estimate of the FP rate of 0.44\\%.  Thus, from the 1500 good single-planet candidates 7 PC+FP systems are predicted.  From the sample of 295 that have two transiting candidates, there are 6 that were identified as a PC+FP combination.  Good agreement is found even in our simplified case.  Of course, one needs to properly account for EBs and the entire range of multi-planet candidates and the multitude of PC and FP mixes that can be produced. These considerations are the basis of Papers I and II, which provide predictions of FPs rates that are verified in this paper.  The combined analysis reported in Paper II and this manuscript validates more than 300 new Kepler multi-planet systems.  Paper II introduces the binary star planet hosts Kepler-132, where one star hosts two transiting planets and its companion hosts one transiting planet and Kepler-296, a pair of small stars with a total of 5 transiting planets, the multi-resonant 4-planet Kepler-223 system and two additional planets in the Kepler-80 = KOI-500 system that includes two 3-body resonances, as well as several high-multiplicity systems, including the new 5-planet systems Kepler-102 = KOI-82, Kepler-169 = KOI-505, Kepler-238 = KOI-834 and Kepler-292 = KOI-1364, three new planets orbiting Kepler-84 = KOI-1589 (bringing the total count to 5) and partial validations of the 5-candidate systems Kepler-122 = KOI-232 (4 planets validated) and Kepler-154 = KOI-435  (2 planets validated).  Hundreds of new planetary systems are announced herein, with special attention given to four new planets with radii roughly twice that of Earth located in or near the habitable zones of their host stars.  This paper is organized as follows: In \\S\\ref{sample} a description of the planetary sample is presented.   The adopted stellar parameters are discussed in \\S\\ref{stellarpars}.  Detailed descriptions of the transit models and lightcurve analysis are described in \\S\\ref{lca}.   The process of identifying FPs can be found in \\S\\ref{fpi}. Subsection \\ref{centroids} covers centroid measurements that provide a clean sample of highly probable transiting multi-planet systems which we demonstrate in \\S\\ref{vps} are genuine extrasolar planets.   We conclude with a discussion of the multi-planet population in \\S\\ref{pop}. ",
        "conclusions": "\\label{pop}     After the application of all the tests listed above, we validate 851 extrasolar planets associated with 340 planetary systems.  These systems are expected to have a FP rate that is significantly less than 1\\% due to the reasons listed in \\S\\ref{vps} and the theoretical framework laid out in Papers I and II.  From this sample, 768 candidates in 306 systems have not been previously validated, but are now extrasolar planets validated above a confidence level of 99\\%.  Thus, we introduce Kepler-100 through 405.  From this population there are 106 new planets that have a radius less than 1.25 \\rearth\\ compared to 16 that have been previously validated.  There are 6 planets with incident solar flux, $S$, less than 1.5 times that of the Earth including four planets: KOI-518.03, 1422.04, 1430.03 and 1596.02 (Kepler-174 d, Kepler-296 f, Kepler-298 d, Kepler-309 c respectively) that are new validations that we discuss below.  Figure \\ref{fig:popplot} plots incident flux versus radius and displays our new validations as filled circles. Multi-planet candidates that pass all tests except our centroid criterion are plotted as open circles and single planet candidates after the exclusion of FPs are plotted as small dots.   The falloff in the number of planets below 1 \\rearth\\ is driven by incompleteness due to insufficient S/N.  The falloff in the number of  planets with $S < 1.5$ is due to decreasing transiting probability and incompleteness to longer period events ($>$150 d)  as our sample is based on Q1-Q8 photometry ($\\sim$ 2 years).  As noted by \\citet{lat11}, there is a lack of hot-jupiters in multi-planet systems verified in Figure \\ref{fig:popplot} as a deficit of planets with $S>200$ and \\rpl\\ $> 7$ \\rearth\\ relative to the single planet population.  The FP tests presented in \\S\\ref{fpi} are not sensitive to most hierarchical blends.  A hierarchical blend is a bound stellar binary with a transiting planet orbiting one of the stellar components.  It is not known if widely separated binaries host planetary systems with orbital planes aligned with the stellar orbital plane.  If the alignment distribution is isotropic, then hierarchal, transiting triples in the multis may be rare, however, if alignment is common, say because of star-planet formation processes that favor aligned systems, then the rate of hierarchal, transiting triples could be much larger.  For an isotropic distribution it was shown in \\S5 of Paper II that 4 or 5 systems are likely to have a planet candidate around each stellar component and it is very unlikely that there is more than 1 hierarchical triple multi (3 bound stars, one hosting a transiting planet and the other 2 in a eclipsing configuration).  If the orbital planes of planets around both components of a stellar binary are aligned then we might expect to find a greater number of blends.  In Appendix A, we develop a synthetic population model to test whether hierarchical blends could contribute a large fraction of the observed multi-planet population by comparing the measured value of \\rhostar\\ from our transit models.  We find that the multi-planet population is not dominated by hierarchical blends, but the strongest constraints come from KECK HIRES observations to search for spectroscopic blends (\\S 8.1 of Paper II).  Together, it appears the rate of hierarchical blends is low.  It is important to note that even if any of our validated planets are found to be orbiting a fainter and bound star, they are still planets; however, the stellar parameters listed in Table \\ref{spars} will need to be revised.  The single and multi-planet populations also appear to have different fractions of planets at longer orbital periods. The relative cumulative distribution of the multiple planet systems overtakes the single planet populations at a period of $\\sim$25 days.  There are 1027 and 897 single and multi-planet candidates with periods between 5 and 150 days.   If we separate these samples into short periods, 5-25 day, and long periods, 25-150 days, we find 689 and 627 short period planets and 338 and 270 long period planets for the single and multi-planet populations respectively.  Thus, 55\\% of the multi-planet sample are found in the short period bin compared to 45\\% for the single planet sample.   Explaining this difference seems counterintuitive, as alignment of orbital planes in multi-planet systems would make it more likely to find longer period planets relative to the single planet population under the assumption that long period planets are equally common in multi and single planet systems.  However, there are strong biases in the detection process that generates the raw KOI list. In particular the candidates are found via different numerical and inspection methods.  For example, there are a number of long period single-transit candidates that were found through identification of a single transit event and then the candidate was continuously monitored for additional transits.  \\subsection{New Planetary Systems with Planets In or Near the Habitable Zone}  We discuss herein, validated multiple planet systems that contain a planet that is in the nominal habitable zone of their star.  The location of the habitable zone depends on stellar luminosity (and the orbital period range also depends on stellar mass), so we introduce only those planets whose host stars have been characterized spectroscopically in this section.  %  As we do not have information regarding either the albedo or atmospheric characteristics of these planets (nor of any moons that they might have), we can only make reasonable estimates of the flux of stellar radiation that they intercept, i.e., the amount of insolation that they receive.  We therefore quote results in terms of the average solar flux intercepted by Earth, $S$, which is generally referred to as the solar constant.  For the purposes of our tabulation, we list objects that intercept flux less than 1.5~$S$ (comparable to the flux Venus received 1 billion years ago, see \\citep{kop13} and references therein). For comparison, Venus receives 1.91~$S$.  We don't specify an outer boundary to the HZ because few of the \\ik planets that we have validated have significantly smaller insolation than does Earth, but note that Mars receives an average of 0.43~$S$. Orbital eccentricity, $e$, which generally is unknown, affects the flux of stellar radiation that a planet receives, but the change that it induces in the annual average insolation  is roughly  quadratic in $e$, and as few of the planets in \\ikt's multis have large eccentricities, the magnitude of the change in mean insolation resulting from planetary eccentricity is likely to be small.  The spectrum of stellar radiation received by a planet also affects atmospheric and surface temperatures \\citep{kas93}, but these variations are small compared to uncertainties in estimated insolation and in atmospheric properties.  Nonetheless, we note the effective temperatures of the stellar hosts to aid in investigations by other researchers.  Only six of the planets that we validate intercept less than 150\\% of the radiation flux encountered by Earth.  Two of these orbit KOI-701 (Kepler-62) and have been analyzed in detail by \\cite{Borucki:2013}.  The transits for the four new planets that receive less than 1.5~$S$, KOI-518.03 (Kepler-174 d), 1422.04 (Kepler-296 f) (see Paper II for more detailed discussion), 1430.03 (Kepler-298 d) and 1596.02 (Kepler-309 c), are shown in Figure \\ref{fig:tplots}.   The planets have nominal radii of 2.19, 1.79, 2.50 and 2.51 \\rearth.  \\subsection{Conclusions}  Our work provides a substantial increase in the number of verified exoplanetary systems and demonstrates the ability of the \\ik mission to probe the statistics of exoplanetary systems with a sample that is relativity clean of FPs.  Both transit models and centroid models are used to characterize the photometric data and various tests were used to identify FPs.  The rate of FPs was found to be low relative to the single transiting planet population, in quantitative agreement with theory (Papers I and II). This result demonstrates that 851 planet candidates in multi-planet systems are valid planets at greater than the 99\\% level.  In Appendix A, the multi-planet population was used to investigate the rate of hierarchical blends in multi-planet systems; while no limits on the occurrence rate can be currently set, it was found that the eccentricity distribution of transiting multi-planet systems found by \\ik is significantly different from the planetary distribution found by RV surveys.  The list of validated planets presented is reliable, but the sample suffers from both incompleteness and strong biases.  Many of the candidates that were not validated in this study are still excellent planetary candidates."
    },
    "1402/1402.4531_arXiv.txt": {
        "abstract": "We present a large-scale, interferometric survey of ammonia (1,1) and (2,2) toward the Galactic Center observed with the Australia Telescope Compact Array (ATCA). The survey covers $\\Delta\\ell\\sim1\\degr$ ($\\sim150$\\,pc at an assumed distance of 8.5\\,kpc) and $\\Delta b\\sim0.2\\degr$ ($\\sim30$\\,pc) which spans the region between the supermassive black hole Sgr\\,A* and the massive star forming region Sgr\\,B2. The resolution is $\\sim 20\\arcsec$ ($\\sim 0.8$\\,pc) and emission at scales $\\gtrsim2\\,\\arcmin$ ($\\gtrsim3.2$\\,pc) is filtered out due to missing interferometric short spacings. Consequently, the data represent the denser, compact clouds and disregards the large scale, diffuse gas. Many of the clumps align with the 100\\,pc dust ring and mostly anti-correlate with $1.2$\\,cm continuum emission. We present a kinetic temperature map of the dense gas. The temperature distribution peaks at $\\sim 38$\\,K with a width at half maximum between $18$\\,K and $61$\\,K (measurements sensitive within $T_{\\rm kin}\\sim10-80$\\,K). Larger clumps are on average warmer than smaller clumps which suggests internal heating sources. Our observations indicate that the circumnuclear disk $\\sim1.5$\\,pc around Sgr\\,A* is supplied with gas by the 20\\,kms\\ molecular cloud. This gas is substantially cooler than gas $\\sim3-15$\\,pc away from Sgr\\,A*. We find a strong temperature gradient across Sgr\\,B2. Ammonia column densities correlate well with SCUBA 850\\,$\\mu$m fluxes, but the relation is shifted from the origin, which may indicate a requirement for a minimum amount of dust to form and shield ammonia. Around the Arches and Quintuplet clusters we find shell morphologies with UV-influenced gas in their centers, followed by ammonia and radio continuum layers.\\\\ ",
        "introduction": "\\label{sec:intro} With a distance of only 8.5\\,kpc \\citep[IAU value used throughout the manuscript, but see][who derive $\\sim 8.0$\\,kpc]{rei09} the Galactic Center is obviously by far the most nearby core of a massive Galaxy. It is dominated by a large Galactic Bar potential \\citep[$\\sim 4.5$\\,kpc half-length][]{cab08} and hosts a $\\sim 4\\times10^{6}$\\,\\msun\\ supermassive black hole \\citep{ghe08,gil09}. The area is at the center of the Galactic gravitational potential and exhibits some of the most complex and energetic processes in our Milky Way. The bar forces gas from the disk to be funneled to the center via $x_1$ and $x_2$ orbits \\citep[e.g.][]{sta04,kim11}. Within the bar potential, most of the molecular gas resides in a $\\sim 500\\,$pc region dubbed the Central Molecular Zone \\citep[CMZ; e.g.][]{mor96,oka98,jon12} and can be considered the region where $x_2$ orbits dominate \\citep[e.g.][]{rod06}. The cloud trajectories in the bar potential are self-intersecting which results in a large number of cloud-cloud collisions. The collisions induce shocks that dominate the physical and chemical properties of the interstellar matter in the region, in particular the molecular gas. The shocks inject turbulence into the gas and are a potential heating source. Turbulence is measured by the increased linewidths of the molecular gas in the Galactic Center region, with values typically of tens of \\kms\\ whereas giant molecular clouds in the disk of the Milky Way exhibit linewidths of only a few \\kms.   Cloud-cloud collisions also influence the properties of star formation in the region. Studies \\citep[e.g.][]{ale07,bar10,lu13} find that the initial mass function in the central parts of the Galactic Center region is skewed toward massive stars more than elsewhere in the Local Universe. In addition, the sheer amount and density of molecular material (about a few per cent, or a few $10^{7}$\\,\\msun, of all molecular gas in the Milky Way is concentrated in about 0.01\\% of the volume) coupled with a star formation rate of $\\sim 0.1$\\,\\sfr\\ \\citep[e.g.][]{cro11} produces a large number of high mass stars. Through stellar evolution, they inject mechanical power of $\\sim 10^{40}$\\,erg\\,s$^{-1}$ in the form of strong stellar winds and supernovae, and return about 0.02\\,\\msun\\ of material back to the ISM \\citep[for an overview, see][]{cro11}. A dramatic effect of these energy sources might be the $\\sim 10\\deg$ sized Fermi bubbles that extend perpendicular to the Galactic disk far into the Galaxy's halo \\citep[e.g.][]{su10,cro11,car13}.  The high density and short lifetimes of the massive stars also produce many compact, energetic objects and the CMZ is under the influence of a large flux of UV and X-ray photons as well as cosmic rays that can penetrate the molecular gas. Some of these sources influence gas locally, others influence gas on larger scales. Overall, however, the state of the gas in the CMZ is dominated by shock physics \\citep[e.g.][]{mar00,mar08,pro08,amo09}.  Given all local and global influences, it is important to understand how star formation can proceed under these conditions. The molecular gas defines the site and the material for star formation and it is an important goal to directly measure the parameters of the molecular gas such as temperature, density, turbulence, ionization fraction and magnetic field penetration. In this paper we concentrate on the temperature of the gas. Our goal is to produce a temperature map of the Galactic Center gas, in particular of the denser gas phase on parsec scales that match the size of the individual star formation locations. A comparison of the temperatures with nearby sources, and a correlation with the cloud properties will allow us to derive the most likely heating sources of the gas, and whether they are external or internal to the gas clouds.  We decide to use ammonia as a temperature tracer. Ammonia (\\amm) has a peculiar structure of a prolate symmetric top, a tetrahedron with a nitrogen atom in one corner and three hydrogen atoms in the remaining corners. The nitrogen atom can tunnel through the plane defined by the three hydrogen atoms. This possibility removes the degeneracy of the symmetric top levels ($J$,$K$) at $J>0$, which splits the energy levels into doublets. Transitions between those levels can be measured as 'inversion lines'. The inversion line frequencies directly depend on the rotational states of the molecule. By measuring the intensity of an inversion line one can thus derive the number of \\amm\\ molecules a the given rotational state. The population interchange between different K ladders is solely governed by collisions, and all ($J>K$,$K$) states decay fast into the metastable $J=K$ levels. The distribution of all ($J$,$J$) states is thus a result of the energetics of the collisions, and is simply described by a Boltzmann distribution at a temperature $T$. Ammonia is thus an excellent thermometer.   In the following we present interferometric observations of the ammonia (1,1) and (2,2) transitions across a large fraction in the CMZ which we use to construct a temperature map of the clumped gas in the region at pc size resolution. The observations and data reduction are described in Sect.\\,\\ref{sec:obs}, followed by our results in Sect.\\,\\ref{sec:results}. A discussion of the morphology, kinematics, temperature structure of the gas clumps and their relation to other Galactic Center objects are presented in Sect.\\,\\ref{sec:discuss}, followed by a summary in Sect.\\,\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Comparison to SCUBA observations} \\label{sec:scuba} In Fig.\\,\\ref{fig:scuba} we show ammonia (1,1) column density contours overlaid on a SCUBA 850\\,$\\mu$m map.  The absolute flux of the SCUBA data is somewhat uncertain due to the chopping technique used \\citep[see][for a description of the original data]{pie00} and should be treated as a lower limit since extended flux is likely filtered out. As mentioned in Sect.\\,\\ref{sec:mos}, our interferometric ammonia measurements suffer from similar problems and a short spacing correction with single dish data would be required to recover the entire flux of the region. Therefore, on large scales, the comparison of extended emission in both maps will not be very reliable. On more compact scales, on which this paper focuses, however, the data can be compared and the ammonia (1,1) column densities and the SCUBA 850\\,$\\mu$m data trace each other extremely well.  All features, except for Sgr\\,A* and the ammonia absorption against the compact \\hii\\ regions Sgr\\,B2(N) and (M) are reproduced in both datasets. Sgr\\,A* is dominated by non--thermal emission \\citep[e.g., ][]{liu01} and is therefore unrelated to dusty molecular gas.  \\placefigure{fig:scuba}  To determine the validity of a dust--to--gas relationship between ammonia and the SCUBA 850\\,$\\mu$m data, we determined both fluxes in each pixel (4\\,\\arcsec\\ in size) after smoothing to a common resolution of 30\\arcsec. Following this, we binned the distribution of the pixels in the \\amm\\ (1,1) column density versus the SCUBA 850\\,$\\mu$m fluxes in order to obtain a point density distribution of these values. The logarithmically scaled point density distribution is shown in Fig.\\,\\ref{fig:dustgas}. The tip of the contours (largest values in dust and gas) follow an almost linear relation as indicated by the dashed line. The line, however, is not centered on the origin of the plot but shifted along the SCUBA 850\\,$\\mu$m axis. This offset may be the result of the missing large scale flux of both maps. However, if the 850\\,$\\mu$m flux is underestimated, the line would shift even further away from the origin on this axis. On the other hand, if it is mainly the interferometric ammonia measurements where we miss flux, the graph would shift up and could be brought in agreement with starting at the origin. But such a shift would be extreme as a correction of $\\sim 2\\times10^{15}$\\,cm$^{-2}$ is required, a value that would almost double the column density for an average cloud (taking account missing SCUBA flux will increase the difference even further). Temperature variations of gas and dust are unlikely to have a significant effect on the slope of the plot, but could broaden the distribution. We also corrected for the temperature and calculated {\\it total} ammonia column densities from the (1,1) line and the temperature \\citep[via the Eq. A15 in][]{ung86}. As listed in Table\\,\\ref{tab:derived}, the (1,1) and (2,2) populations are together about half of the total ammonia column density that we derive from our data.  A comparison with the entire ammonia column density thus leads to a very similar plot as the one shown in Fig.\\,\\ref{fig:dustgas} and it does not contain a substantial additive component. In addition, 850\\,$\\mu$m dust fluxes are not exceedingly susceptible to variations in temperature. The offset of the line may thus be due to a large amount of missing flux from extended emission mainly of the \\amm(1,1) line. A different explanation for the offset would be that the ammonia forms on dust grains which also shields the material and that a minimum of dust is required to produce ammonia efficiently and to protect it against destruction through the surrounding UV radiation field \\citep[see also the discussion in][]{wei01}.   \\placefigure{fig:dustgas}   \\subsection{The Radio Filaments and the Molecular Gas} \\label{sec:filaments} In contrast to the SCUBA maps, the radio continuum emission and the molecular gas as traced by ammonia show little coincidence. Only the star forming region Sgr\\,B2 exhibits radio continuum emission from the \\hii\\ regions at the same positions where large amounts of molecular gas are found (see Fig.\\,\\ref{fig:m0}). For this region it is obvious that the molecular gas most likely feeds the current star forming activity. As already noted by \\citet{meh95}, there is much less molecular gas in the star forming Sgr\\,B1 region at $(\\ell,b)=(0.50\\degr,-0.05\\degr)$. This may indicate that the Sgr\\,B1 region is more evolved than Sgr\\,B2 and that the remaining molecular gas has already been dispersed.  In addition, the non-thermal radio arcs at $\\ell\\sim0.15$ also occupy a region with only little ammonia line emission. The pistol/sickle region at $(0.16\\degr,-0.06\\degr)$, for example, exhibits \\amm(1,1) column densities of only $\\sim 2\\times10^{13}$\\,\\cden. In addition, at larger Galactic latitudes, the radio arc bifurcates into two roughly parallel components at $\\ell\\sim 0.19\\degr$ and $\\ell\\sim0.17\\degr$. Along the filaments, the \\amm(1,1) column densities are $\\sim 12\\times10^{14}$\\,\\cden, whereas the column density in the very thin, long area in between the two filaments rises up to $\\sim 28\\times 10^{13}$\\,\\cden. The radio emission in the two arched thermal filaments that connect to the main non-thermal arcs and that stretch from $(0.14\\degr,0.05\\degr)$ to $(0.08\\degr,0.00\\degr)$ and from $(0.10\\degr,0.08\\degr)$ to $(0.06\\degr,0.02\\degr)$ are also in regions with low \\amm\\ average column densities. Toward lower Galactic latitude and lower Galactic longitudes, however, very prominent molecular gas clumps are apparent in the integrated intensity maps (c.f. Fig.\\,\\ref{fig:m0}). A notable exception to this behavior is the ``handle'' of the pistol/sickle, the radio filament that extends roughly perpendicular to the large non-thermal radio arcs at Galactic latitudes of $\\sim -0.04\\degr$. Molecular gas appears to be aligned with this feature and exhibits column densities of $\\sim 2.4\\times10^{14}$\\,\\cden. Finally, the region around Sgr\\,A* with the associated mini--spiral is also very bright in 1.2\\,cm radio continuum emission but has only weak features of molecular gas. In particular, the very prominent circumnuclear, molecular ring around Sgr\\,A*, which is very bright in HCN \\citep{chr05}, is not exceptionally strong in the \\amm (1,1) and (2,2) transitions (see Sect.\\,\\ref{sec:cnd}). The ammonia feature has a somewhat larger diameter than that of HCN and column densities of $2-4\\times 10^{14}$\\,\\cden. This is up to an order of magnitude less than the column densities in the 20 and 50\\,\\kms\\ clouds. The more isolated radio continuum regions at $(-0.013\\degr,0.021\\degr)$ and $(-0.053\\degr,0.020\\degr)$ also show local peaks in the ammonia distribution of a few $10^{14}$\\,\\cden. Associated star formation may thus be in a relatively early stage.     \\subsection{The Molecular Arcs and the 100\\,pc Dust Ring} \\label{sec:arcs}   Some of the most prominent features in the interferometric ammonia map (Fig.\\,\\ref{fig:m0}) are the two molecular arcs west of Sgr B2, between $\\ell\\approx0.2$\\degr\\ and $0.6$\\degr. The two Galactic northern and southern arcs are almost symmetric in shape, mirrored along a line at constant $b\\approx-0.02$\\degr. The northern arc is certainly much brighter than the low-b counterpart. In Fig.\\,\\ref{fig:scuba}(c) we show an overlay of the ammonia to {\\it Spitzer} 8\\,$\\mu$m data. Some of the detected gas corresponds very well to strong absorption features in the infrared images. In particular the northern arc is heavily absorbing the infrared emission whereas the southern arc is hardly visible in the {\\it Spitzer} image at all. Similarly, the string of molecular clouds north of Sgr A* are less absorbing than those south of Sgr\\,A*, namely the 20 and 50\\,\\kms\\ clouds. This can be interpreted such that the northern arc and the clouds south of Sgr\\,A* are in front of most of the IR emission whereas the southern arc and the clouds north of Sgr\\,A* are located behind. The different structures are indicated by the sinusoidal solid (back structure) and dashed (front) lines in Fig.\\,\\ref{fig:vel}(b) (also in Fig.\\,\\ref{fig:pv}).  This picture is consistent with the morphology of the ``figure 8'' 100\\,pc ring that was recently detected in a {\\it Herschel} temperature map by \\citet{mol11}. The shapes of the dust ring and our ammonia arcs are very similar and the ammonia may thus be emitted by clumps residing in the dust ring.  \\citet{mol11} claim that the ring is rotating, based on single dish data of the CS molecule. Our data is filtering out the diffuse emission and the clumps seem to follow a somewhat different picture. Whereas the \\citet{mol11} interpretation suggests that the foreground arc east of Sgr A* should be redshifted compared to the background arc, we find the opposite. In Fig.\\,\\ref{fig:pv} the northern, foreground arc (dashed line), is at lower velocities compared the foreground arc (solid line). The velocities of the two arcs are less distinct toward Sgr B2 and at the position of the non-thermal arcs at about $b\\sim0.15\\degr$. This indicates that the structure might be described by an expansion with a velocity of $\\sim 30$\\,\\kms. Signatures for an expansion were also detected in previous CO measurements by \\citep[][]{sof95}, albeit at much larger expansion velocities. A comparison of our ammonia data with the intensity-weighted N$_2$H$^+$ map presented in \\citet{jon12} corroborates the view gained on the basis of the ammonia data alone, i.e. it appears to contradict the rotating ring picture. Our map, unfortunately, does not cover the region between Sgr\\,A* and Sgr\\,C, which contains the other half of the 100\\,pc ring at $\\ell<0\\degr$.   \\subsection{The Temperature Distribution of Molecular Clumps} \\label{sec:temp}  Due to the filtering effect of the ATCA interferometer, we derive the temperature distribution of molecular clumps more or less directly, with little contribution from diffuse inter-clump gas.  In Fig.\\,\\ref{fig:tkinnh}(a), we show the kinetic temperature distribution of all pixels in our map. The distribution peaks at about 38\\,K and has a FWHM of 43\\,K, spanning the 18-61\\,K range. We derive very few areas with a lower limit of kinetic temperatures of $>80$\\,K. Previous studies have detected high temperature components in the Galactic Center molecular gas using single dish telescopes \\citep[e.g.][]{gue81,hue93b,rod01,ao13,mil13}. It is very likely that the hot material is mostly diffuse and filtered out by our interferometric observations in combination with that our (1,1) and (2,2) lines are not sensitive to high temperatures. This pronounces the clumpier material which is substantially colder on average.  \\placefigure{fig:tkinnh}  From Fig.\\,\\ref{fig:tkin}(b) we may see a trend that smaller clumps exhibit colder temperatures. A different view is shown in Fig.\\,\\ref{fig:tkinnh} where we plot the logarithm of the total ammonia column density (cf. Sect.\\,\\ref{sec:scuba}) to the derived kinetic temperature of the gas. The pixel by pixel correlation has been binned and logarithmic contours are shown. The plot confirms that gas at higher column density also tends to show a higher temperature. If the heating is external, one would expect it to affect clouds of all sizes to similar depths. Smaller clumps would thus be more thoroughly heated whereas larger clumps would be able to maintain a colder inner core. Since we see the opposite, we may assume that most of the gas heating takes place internal to the individual clumps. If the star formation rate per molecular gas mass is roughly constant, larger clumps will contain more proto-stars, in particular more massive proto-stars. Those stars are the most likely candidates for internal heating. We hardly detect massive clumps with very low temperatures, or small clumps with large temperatures. Small clumps that do contain massive proto-stars may in fact dissipate relatively quickly, eluding their observation.   The Galactic Center, however, is known to contain external heating sources such as cosmic rays, shocks, a strong UV field, X-ray sources, and dust. As explained above, most of the heating in the clumps is likely internal. But in some places the external heating may play a significant role, even for the smaller clumps. Notably, this includes the bow-shaped region from the thermal, arched filaments, through the high column density 50\\,\\kms\\ and 20\\,\\kms\\ clouds, in addition to the compact and dense Sgr B\\,2 region which exhibits a strong temperature gradient (both traced by the yellow and red spots in Fig.\\,\\ref{fig:tkin}[a]). The bow-shaped region is close to Sgr\\,A*. Although the CND is cooler (see Sect.\\,\\ref{sec:cnd}), the warmer gas may be influenced by activity of Sgr\\,A*, which may cause larger X-ray and cosmic ray fluxes, as well as shocks. PDR effects may play a major role in the ``mini-starburst'' Sgr\\,B2, which is also a place where the $x_1$ and $x_2$ orbits of the central stellar bar intersect. The resulting shocks might be an additional heating mechanism for the molecular gas in that area and can potentially explain the Sgr\\,B2 temperature gradient that is hotter at larger Galactic longitude values.    \\subsection{The Arches and Quintuplet Cluster Regions} \\label{sec:clusters}  Besides the Sgr\\,B regions, young stellar clusters are visible all across the CMZ. Two of the most prominent, massive clusters are the Arches $(0.121\\degr, 0.017\\degr)$ and the Quintuplet $(0.160\\degr, -0.059\\degr)$ clusters displayed in Fig.\\,\\ref{fig:clusters}. The clusters have masses of $\\sim 10^{4}$\\,\\msun\\ and are with ages of $\\sim 2$\\,Myr for the Arches and $\\sim 4$\\,Myr for the Quintuple cluster \\citep{fig99} relatively young. The different ages may also be reflected in the dense gas map. Close to the Arches cluster, a large filament of dense molecular gas extends roughly parallel to the non--thermal radio continuum filament. The immediate surroundings of the Quintuplet cluster, however, are relatively clear of ammonia and may have dispersed over the lifetime of this cluster. The three--color composite map of the region in Fig.\\,\\ref{fig:clusters} shows how different the distributions of the gas tracers are. CO emission \\citep[][]{mar04} is brightest closer in the direction to the Galactic Center and only little is visible in a roughly circular region that is defined by the radio arcs, the pistol/sickle and the thermal arched radio filaments that extend from the non-thermal arc to smaller Galactic longitudes. This region can be roughly described by a circle centered on $(\\ell,b)=(0.138\\degr,-0.007\\degr)$ and having a diameter of $0.08\\degr$ (indicated by the cyan circle in Fig.\\,\\ref{fig:vel}[b]). The region, almost entirely surrounded by radio continuum emission appears to be filled with [\\ion{C}{1}] $^{3}P_{\\rm 0}-^{3}P_{\\rm 1}$ emitting gas \\citep[rest frequency: 492.161\\,GHz; data from][]{mar04}. [\\ion{C}{1}] is a well--known tracer for photo--dissociation regions (PDR) and this indicates that the gas is mainly ionized within this cavity. This is supported by the detection of Fe\\,K$\\alpha$ 6.4\\,keV emission in the region \\citep{wan02,wan06,tat12}. The gas traced by ammonia is scattered all across that area. The main filaments, however, are close to the radio continuum emission toward the inside of the region. The two stellar clusters are also at the inner edge of the radio continuum emission. The ammonia channel maps (Fig.\\,\\ref{fig:chann11}) at this position show signatures that may be interpreted as an expanding structure but the relative velocities could well be due to other motions, too.  If it is expanding, the driving force could be provided by the gas traced by [\\ion{C}{1}]. This scenario would also explain why the morphology and kinematics of the \\amm\\ emission does not follow the overall sinusoidal pattern that is prominently observed by the molecular arcs between the shell (cyan circle in Fig.\\,\\ref{fig:vel}[b]) and Sgr\\,B2, which is continued south and north of Sgr\\,A* (Sect.\\,\\ref{sec:arcs}).  \\placefigure{fig:clusters}  The presence of strong [\\ion{C}{1}] emission also indicates PDR conditions and hence a higher ionizing flux and, since the [\\ion{C}{1}] emission is widespread, a higher temperature. The temperature map of ammonia (Fig.\\,\\ref{fig:tkin}), however, shows values of $T_{\\rm kin}\\sim 40-60$\\,K which are only slightly enhanced compared to other, more quiet regions of the mapped area. The denser gas therefore appears to be well--shielded against heating from ionizing flux. If the picture of an expanding shell structure holds true, the distribution of the different gas tracers may thus indicate that dense gas forms a boundary layer that is filled with ionized and PDR gas (see also the discussion in Sect.\\,\\ref{sec:filaments}). The dense gas layer is also where the Arches and Quintuplet clusters are located. The radio continuum is emitted in an additional sheet surrounding the dense gas that takes the role of the interface to the unaffected ISM. Substructures are observed close to the stellar clusters and within the radio continuum emission.  The more immediate environment of the Quintuplet cluster is relatively devoid of ammonia emission but has a local peak in [\\ion{C}{1}]. It also shows radio continuum emission aligned with the outer part of the structure. This emission forms the pistol/sickle region. The situation is less clear for the gas in the immediate surroundings of the Arches cluster. The Arches appears to be located on the rim of a cavity, that is centered on $(\\ell,b)\\sim(0.107\\degr,0.033\\degr)$ and filled with PDR gas. Since warm, ionized gas follows the path of least resistance, it may originate from the Arches cluster but then accumulate in this substructure.     \\subsection{The Circumnuclear Disk} \\label{sec:cnd}  The circumnuclear disk around Sgr\\,A* is the most central position in our Milky Way at which substantial amounts of molecular gas are detected. The ring has been mapped in various dense molecular tracers, like HCN \\citep[e.g.][]{chr05,mon09,mar12} and shows a characteristic diameter of an arcminute (or about 3\\,pc) with a total mass of $\\sim 10^{6}$\\,\\msun\\ stored in various clumps \\citep{chr05}. A temperature map with \\amm (1,1) column density contours is shown in Fig.\\,\\ref{fig:cnd}. We detect a clumpy structure that is in broad agreement with the HCN observations. \\citet{her05} mapped the region in the ammonia lines with the Very Large Array. They detect an even clumpier distribution with their more extended array. With our shorter spacings, we are able to detect the most likely connection of the CND with the 20\\,\\kms\\ cloud and possibly the 50\\,\\kms\\ cloud \\citep[see also][]{mcg01}. The 20\\,\\kms\\ cloud has a velocity gradient and at the position where it connects to the CND, the value is about $40$\\,\\kms\\ (see Fig.\\,\\ref{fig:cnd}). From this level, the values roughly decrease in a clockwise direction around the CND. This indicates that gas is fed from the 20\\,\\kms\\ to the CND.  Given that the CND is only $\\sim 1.5$\\,pc away from the supermassive black hole Sgr\\,A*, it is surprisingly cold. We detect kinetic temperatures between 25\\,K in the clumps and a maximum of $\\sim 70$\\,K in some inter-clump regions. Our values are much lower than single dish measurements that derive a few hundred K for two regions of the CND \\citep[e.g.][]{req12}. This could indicate that the CND is composed of clumpy cold and diffuse hot material as we are not sensitive to extended emission.  The temperature structure of the clumps in the CND is not significantly different from regions further away, e.g. in the molecular arcs (Sect.\\,\\ref{sec:arcs}). As seen in Fig.\\,\\ref{fig:tkin}, the temperatures of the CND are also lower than those in a the warmer bow-shaped region described in Sect.\\,\\ref{sec:temp}. Cooling and shielding appears to be more efficient in the CND, if the higher surrounding temperatures are caused by Sgr\\,A*.  \\placefigure{fig:cnd}      We observed the region between the supermassive black hole Sgr\\,A* and the massive star forming region Sgr\\,B2 in our Milky Way with the ATCA in the ammonia (1,1) and (2,2) inversion lines ($-0.2\\degr \\leqslant l \\leqslant 0.8\\degr$; $-0.13\\degr \\leqslant b \\leqslant 0.07\\degr$). The interferometric data filters the clumpy structure that otherwise blends with a more diffuse component in single dish observations. We find the following:   \\begin{itemize}  \\item The dense gas is distributed in many filamentary structures with two prominent molecular arcs between Sgr\\,A* and Sgr\\,B2. The dense molecular arcs are likely located on the '100 pc' ring that was recently detected in {\\it Herschel} dust temperature maps.  \\item Except for the extreme Sgr\\,B2 region, the 1.2\\,cm radio continuum and ammonia emission are almost anti-coincident. This includes the continuum emitted from Sgr\\,B1, the non-thermal arcs, the arched, thermal filaments, as well as the mini-spiral around Sgr\\,A*.  \\item SCUBA 850$\\mu m$ data trace the dense ammonia gas very well, in an almost linear relationship. We see possible indications that a minimum amount of dust may be required for substantial amounts of ammonia to form and survive the encompassing UV radiation field.   \\item A map of kinetic temperatures is constructed that is reasonably accurate in the 10-80\\,K range. For some of the gas, in particular gas close to Sgr\\,B2 and in the vicinity of the 20 and 50\\,\\kms\\ clouds and the arched, thermal filaments, we derive a lower limit of $>80$\\,K for the kinetic temperature.   \\item The peak of the kinetic temperature distribution is at $\\sim 38$\\,K with a FWHM that spans 18-61\\,K. On average, larger clumps exhibit larger temperatures than smaller clumps. The dominating heating mechanism is thus likely internal, and possible sources could be proto-stars.  \\item The molecular gas for the circumnuclear disk immediately surrounding Sgr\\,A* with a radius of $\\sim 1.5$\\,pc is likely being supplied by the 20\\,\\kms\\ cloud. The CND is cooler than gas in the 20\\,\\kms\\ and 50\\,\\kms\\ clouds, and gas closer to the arched, thermal filaments.    \\item The region close to the Arches and Quintuplet clusters shows signs for a layered structure with more diffuse PDR gas that is surrounded by dense ammonia clumps. Furthermore, the radio continuum filaments form an outer layer.  \\end{itemize}  The Galactic Center is an energetic region and understanding the dense gas clumps is key to understanding star formation in extreme conditions. The difference of the diffuse and dense gas components can be revealed by additional single dish observations. Finally it will be desirable to observe higher transitions over the full CMZ to obtain a map of a larger range of temperatures. For example, the \\amm(2,2) line has an energy of $\\sim 64$\\,K above ground, but the (3,3), (6,6) and (9,9) levels are at $\\sim 123$\\,K, $\\sim 407$\\,K, and $\\sim 852$\\,K, respectively, and therefore trace much warmer gas. Other temperature tracers can be used to calibrate all methods against each other. The CMZ is also an important benchmark for comparisons with other galactic nuclei, near and far, active and quiet. Understanding the details of heating and cooling, star formation, AGN feedback and gas transfer in the nearby CMZ is therefore indispensable to understand the role of nuclear regions for galaxy evolution at all cosmic epochs."
    },
    "1402/1402.6702_arXiv.txt": {
        "abstract": "We study how the estimation of the sonic Mach number (\\ms) from $^{13}$CO linewidths relates to the actual 3D sonic Mach number. For this purpose we analyze MHD simulations which include post-processing to take radiative transfer effects into account. As expected, we find very good agreement between the linewidth estimated sonic Mach number and the actual sonic Mach number of the simulations for optically thin tracers.  However, we find that opacity broadening causes \\ms ~to be overestimated by a factor of $\\approx1.16-1.3$ when calculated from optically thick $^{13}$CO lines. We also find that there is a dependency on the magnetic field: super-Alfv\\'enic turbulence shows increased line broadening as compared with sub-Alfv\\'enic turbulence for all values of optical depth for supersonic turbulence.  Our results have implications for the  observationally derived sonic Mach number--density standard deviation ($\\sigma_{\\rho/<\\rho>}$) relationship, $\\sigma^2_{\\rho/<\\rho>}=b^2\\mathrm{M_s^2}$, and the related column density  standard deviation (\\sign) sonic Mach number relationship. In particular, we find that the parameter $b$, as an indicator of solenoidal vs. compressive driving, will be underestimated as a result of opacity broadening. We compare the \\sign--\\ms~relation derived from synthetic dust extinction maps and $^{13}$CO linewidths with recent observational studies and find that solenoidally driven MHD turbulence simulations have values of \\sign ~which are lower than real molecular clouds.  This may be due to the influence of self-gravity which should be included in simulations of molecular cloud dynamics. ",
        "introduction": "Supersonic magnetized turbulence is observed in multiple tracers across several different interstellar media (ISM) phases. This includes the neutral medium, as traced by HI \\citep[see][]{2010ApJ...714.1398C,2010ApJ...708.1204B,2011ApJ...735..129P}, the warm ionized medium, as traced by H$\\alpha$ \\citep[see][]{2008ApJ...686..363H}, electron density fluctuations \\citep{1995ApJ...443..209A,2010ApJ...710..853C}, synchrotron polarization \\citep{2011Natur.478..214G,2012ApJ...749..145B} and the molecular medium \\citep[see][]{1998ApJS..115..241H,2008ApJ...680..428G}, which includes a variety of molecular traces including the often used carbon monoxide (CO) line.  Turbulence and magnetic fields in these environments are interrelated to a variety of physical process, including cosmic ray transport \\citep{2004ApJ...614..757Y,2011ApJ...728...60B}, magnetic reconnection \\citep{1999ApJ...517..700L,2009ApJ...700...63K}, and star formation \\citep[and ref. therein]{2007ARA&A..45..565M}.  Furthermore it is clear that, in order to understand MHD turbulence and related physical mechanisms, one needs to be able to measure basic plasma parameters such as the sonic and Alfv\\'enic Mach numbers ($M_s\\equiv V_{turb}/c_s$, where $c_s$ is the sound speed, and $M_A\\equiv V_{turb}/V_{Alfven}$, respectively).   It is clear from simulations, observations, and theoretical works that compressible turbulence, which generates shocks, is important for creating filaments and local regions of high density contrast \\citep{2007ApJ...666L..69K}. Shocks broaden the gas/dust density and column density probability distribution function (PDF) as well as increase the peaks and drag out the distribution's tails towards higher gas densities \\citep{2009ApJ...693..250B}. Based on this observation, several authors \\citep[e.g.][to name a few]{1994ApJ...423..681V,1997ApJ...474..730P,2010A&A...512A..81F,2012ApJ...755L..19B} have developed relationships between the PDF moments, such as the variance or standard deviation of density ($\\sigma_{\\rho/<\\rho>}$), and the sonic Mach number, for example  \\begin{equation} \\sigma^2_{\\rho/<\\rho>}=b^2M_s^2\\label{sigbm} \\end{equation}  where b is a parameter that depends on the type of turbulence forcing with b=1/3 for pure solenoid forcing and b=1 for pure compressive forcing \\citep{2008ApJ...688L..79F}. Once gas becomes dense enough, collapse can occur via self-gravity and the PDF forms power-law tails towards higher density regions \\citep{2000ApJ...535..887K, 2012ApJ...750...13C}. Additional variants of Equation 1 have also be developed, e.g. including plasma $\\beta$ \\citep[see][]{2012MNRAS.423.2680M}.  Eq. \\ref{sigbm} is hard to constrain observationally as volume densities and the 3D velocity structure are not available from observations. Important information on turbulent supersonic motions in molecular clouds comes from the observed non-thermal broadening of the linewidths of different spectral lines, e.g. from carbon monoxide (CO) emission.  However, CO, in particular \\co ~is often partially or fully optically thick and traces only a limited dynamic range of column densities \\citep[see][b]{2009Natur.457...63G,2013ApJ...771..122B}.   Fortunately, dust extinction column density maps of infrared dark clouds (IRDCs), including mid infrared (MIR) and near infrared (NIR) wavelengths can be used to trace a much larger dynamic range of densities in order to probe the PDF \\citep[Av=1--25 for NIR and Av=10--100 for MIR, see][henceforth KT13]{2001A&A...377.1023L,2011A&A...530A..64K,2013A&A...549A..53K}. However, extinction maps carry no dynamical information regarding velocities and for this, molecular line profiles are needed to measure the dynamics of clouds, including the sonic Mach number for a given cloud temperature.  In this paper we investigate the robustness of measuring the sonic Mach number from maps of $^{13}$CO using synthetic observations derived from 3D MHD turbulence simulations.  We further create synthetic dust maps in order to compare the PDF standard deviation with the measured sonic Mach number, following the approach in the observational work of KT13. ",
        "conclusions": "We create synthetic \\co ~emission maps, with varying optical depth,  and dust column density maps from a set of 3D MHD simulations. We derive a standard deviation - sonic Mach number relation, as it would be found from the observations, and compare this with recent results from real clouds discussed in KT13.  We find:  \\begin{itemize} \\item  Calculations of \\ms ~from line widths of \\co ~are robust for optically thin \\co ~but are overestimated by a factor of up to $\\sim1.3$ for optically thick clouds.  This is due to the well known, but often overlooked, effect of opacity broadening. \\item This over estimation of the sonic Mach number as derived from \\co ~line widths will cause the slope of the \\sign-\\ms ~relation to become more shallow and this will result in lower values of the measured b parameter. \\item The \\sign~values of clouds reported in KT13 are larger than values found from ideal solenoidally driven simulations of turbulence.  This could be due to the fact that in real molecular clouds there exists the influences of gravity, which will increase the measured column density standard deviation. \\end{itemize}"
    },
    "1402/1402.5747.txt": {
        "abstract": "We carried out the photometric observations of the SU UMa-type dwarf nova ER UMa during 2011 and 2012, which showed the existence of persistent negative superhumps even during the superoutburst. We performed two-dimensional period analysis of its light curves by using a method called ``least absolute shrinkage and selection operator'' (Lasso) and ``phase dispersion minimization'' (PDM) analysis, and we found that the period of negative superhumps systematically changed between a superoutburst and the next superoutburst.  The trend of the period change can be interpreted as reflecting the change of the disk radius. This change of the disk radius is in good agreement with the predicted change of the disk radius by the thermal-tidal instability (TTI) model.  The normal outbursts within a supercycle showed a general trend that the rising rate to maximum becomes slower as the next superoutburst approaches.  The change can be interpreted as the consequence of the increased gas-stream flow onto the inner region of the disk as the result of the tilted disk. Some of the superoutbursts were found to be triggered by a precursor normal outburst when the positive superhumps appeared to develop.  The positive and negative superhumps co-existed during the superoutburst. The positive superhumps were prominent only during four or five days after the supermaximum, while the signal of the negative superhumps became strong after the middle phase of the superoutburst plateau.  A simple combination of the positive and negative superhumps was found to be insufficient in reproducing the complex profile variation. We were able to detect the developing phase of positive superhumps (stage A superhumps) for the first time in ER UMa-type dwarf novae.  Using the period of stage A superhumps, we obtained a mass ratio of 0.100(15), which indicates that ER UMa is on the ordinary evolutional track of cataclysmic variable stars. ",
        "introduction": "Dwarf novae (DNe) are a class of cataclysmic variables (CVs), which consist of a white dwarf primary and a late-type secondary which fills its Roche lobe. The material transfered toward the primary through the inner Lagrangian point (L1) forms an accretion disk around the white-dwarf. The accretion disk causes instabilities, which are observed as an outburst [for reviews, see \\citet{war95book}; \\citet{osa96review}; \\citet{hel01book}]. SU UMa-type stars are a subgroup of dwarf novae. They are characterized by the presence of two types of outbursts, normal outburst and superoutburst. Whereas a normal outburst lasts for only a few days, a superoutburst lasts for about two weeks and the maximum magnitude of the latter is brighter by 0.5--1 mag.  These objects also exhibit light variations called (positive) superhumps during superoutburst. The observed period of the superhumps is a few percent longer than the orbital period of the system. The positive superhumps are thought to arise by periodic viscous dissipation of tidally elongated disk (i.e. the eccentric disk) whose aspsidal line slowly precesses in the prograde direction (see \\cite{whi88tidal}, \\cite{hir90SHexcess} ) % This phenomenon is well explained by a dynamical prograde %precession of the elongated disk. On the other hand, some cataclysmic variables show variations shorter than the orbital period called ``negative superhumps'' (\\cite{uda88ttari}; \\cite{har95v503cyg}; \\cite{rin12v378peg}). The origin of negative superhumps is usually considered as a result of retrograde precession in a tilted accretion disk \\citep{woo07negSH}.  When the disk is tilted, the hot spot is formed at the inner part of the disk, not the edge of the disk. Since the energy of hot spot comes from the release of gravitational energy, such change in the location of the hot spot causes a variation in the amount of the released energy, namely the luminosity of the hot spot. Combined with the retrograde precession, this effect can explain the negative superhumps. %\\citet{lub92tilt} proposed the tidally inclination instability %as the source of disk tilts. %\\citet{mon10disktilt} recently proposed hydrodynamic lift %as the common source of disk tilts, and indicated that %there is a minimum mass-transfer rate to generate tilts. %This suggestion appears to be consistent with the observation %that negative superhumps are mainly detected in %systems with high mass-transfer rate.  The interval of successive two superoutbursts is called ``supercycle''. These two superoutbursts usually sandwich several normal outbursts. In order to explain such behavior of SU UMa-type dwarf novae, the tidal thermal instability (TTI) model is suggested \\citep{osa89suuma}. In this model, systems with small mass ratios ($M_{2}/M_{1} = q \\le 0.25$) enable the disk to reach the radius of the 3:1 resonance to the orbital motion of the secondary. In normal outbursts, the material of the disk only partly accrete to the inner region. The radius of the disk becomes gradually larger after experiencing normal outburst. When the disk radius reaches the 3:1 resonance radius, the eccentric instability is excited. The increased turbulence in the disk causes an increase in the mass-transfer rate in the disk and a long, bright superoutburst is triggered \\citep{osa89suuma}. This prograde precession also causes the superhump (\\cite{whi88tidal}, \\cite{hir90SHexcess}).  ER UMa is a member of SU UMa-type dwarf novae and its intervals of superoutburst (supercycle) are as short as 40 -- 50 d \\citep{kat95eruma}. This object is known as the prototype of a subgroup, ``ER UMa type'' which is characterized by having extremely short supercycles ($<$ 60 d) among SU UMa-type stars (e.g. \\cite{rob95eruma}, \\cite{nog95v1159ori}; for a review see \\cite{kat99erumareview}). Although \\citet{gao99erumaSH} and \\citet{kju10eruma} suggested on the presence of negative superhumps during quiescence and a normal outburst, only positive superhumps were observed during the following superoutburst. However, \\citet{ohs12eruma} reported that negative superhumps were detected in ER UMa during the superoutburst in 2011 January. This is the first confident detection of negative superhumps during the superoutburst of any SU UMa-type dwarf nova. In \\citet{ohs12eruma} (hereafter paper I), we reported the persistent negative superhump was detected in ER UMa and implied the possibility that the existence of negative superhumps suppresses of the occurrence of normal outbursts .  The existence of negative superhumps during superoutburst was also reported in other SU UMa-type dwarf novae V1504 Cyg and V344 Lyr (\\cite{osa13v1504cygKepler}, \\cite{sti10v344lyr}). \\citet{osa13v1504cygKepler} analyzed data of an SU UMa-type dwarf nova V1504 Cyg observed by Kepler, and showed that this object also shows negative superhumps during superoutburst as well as during normal outbursts and in quiescence. \\citet{osa13v1504cygKepler} have demonstrated that in V1504 Cyg the frequency of occurrence of normal outbursts in a supercycle is related to the existence or non-existence of negative superhumps in the sense that it is reduced when the negative superhumps exist, confirming the suggestion made in Paper I. That is, the cycle lengths of normal outbursts are longer in a supercycle with negative superhumps than those in a supercycle without negative superhumps, and they called the former case ``type L-supercycle'' and the latter case ``type S-supercycle'' (these two types were first introduced by \\citet{sma85vwhyi} for supercycles observed in VW Hyi and symbols ``L'' and ``S'' come from ``long'' and ``short'' for normal outburst cycles).  In Paper I, we dealt with  only one supercycle  of ER UMa in 2011. We have made further comprehensive observations of ER UMa in 2011 and 2012 and, our dat covered three supercycles in 2011 and three supercycles in 2012. All of  supercycles observed have turned out to be were all accompanied with negative superhumps, and provide an excellent opportunity to study the  outburst behavior when negative and positive superhumps co-exist. In this paper, we report on these new observations together with a more sophisticated analysis of the data reported in paper I. We explain our  observations in section \\ref{obschap} and we present the result of their analysis observation and analysis in section \\ref{reschap}. The conclusion is given in section \\ref{conchap}. ",
        "conclusions": "\\label{conchap}  We observed the SU UMa-type dwarf nova ER UMa in the years 2011 and 2012 by VSNET worldwide campaign and we obtained data of 307 nights spaning in 2011 -- 2012. We detected  persistent negative superhumps during the superoutbursts as well as during normal outbursts and quiescence  in both seasons. We analyzed these data and obtained the following are our major findings:  (1) We succeeded in covering three supercycles of this star in 2011 and also three supercycles in 2012. The star showed persisted superhumps in both seasons.  (2) We analyzed periodic variations of negative superhumps in a supercycle of this star by  $O-C$ analysis, PDM analysis, and Lasso analysis.  We have found from these analyses that the period of negative superhumps between two successive superoutbursts decreases secularly from the end of a superoutburst to the next superoutburst in a longer time-scale of a supercycle. Superimposed on this long time-scale trend,  a shorter time-scale variation in the negative superhump period occurred within a normal outburst cycle in a sense that the period was shortened when a normal outburst occurred and it became longer during quiescence.  When the next supercycle started, the period of negative superhumps returned to the value when the previous supercycle started. Thus it varied cyclically with the same supercycle period as that of the light curve. This variation was indicates as a result of the variation in the disk radius and the disk radius was the smallest at the end of a superoutburst. The variation in the disk radius  inferred by our analysis of ER UMa is in agreement with the predicted variation of the disk radius by the TTI theory.  (3) The rising rate of maximum of the normal outburst varied during one supercycle. The rising rate  of the first normal outburst within a supercycle was found to be faster than the later ones. This suggests that the first normal outburst may most likely be of the ``outside-in'' type while the later ones may be of the ``inside-out'' type. The occurrence of the ``inside-out''-type outbursts in ER UMa may be understood as due to the mass supply from the secondary to the inner part of the disk when the disk tilted.  (4) The negative superhumps and the positive superhumps co-existed during the superoutburst although the signal of the negative superhump was marginal in the first several days of the superoutburst. By subtracting the signal of negative superhumps, we obtained the $O-C$ diagram for the positive superhumps and by doing so we were able to distinguish the stage A from the B ones. A simple combination of the positive and negative superhumps was not sufficient to reproduce the complex profile variation.  (5) The number of normal outbursts in a supercycle of ER UMa in 2011 and 2012 was found to be mostly four, which  was smaller than that in other occasions when negative superhumps were not observed,  a result consistent with those reported for ER UMa by \\citet{zem13eruma} and for V1504 Cyg by \\citet{osa13v1504cygKepler}. This can be understood as a result of the tilted disk: when the disk is tilted, the gas stream coming from the secondary arrives at the inner region of the disk rather than at the disk edge and this reduces a frequent occurrence of the outside-in type normal outbursts.  (6) Some of the superoutbursts were triggered by the normal outbursts ( i.e. precursor outbursts). Positive superhumps started to grow during the declining part of the precursor in such superoutbursts. The phases of maxima were found to be continuous when a transition from the negative superhump to the positive superhump occurred.  (6) We succeeded in detecting stage A superhumps for the first time in ER UMa-type dwarf novae during the precursor parts of three superoutburst. Using the period of stage A superhumps, we estimated $q$ of ER UMa to be 0.100(15).  The estimated $q$ and the known orbital period imply that ER UMa itself is in the standard evolutionary track of cataclysmic variable stars. However, the mass transfer late, $\\dot{M}$ of ER UMa,  inferred from its short supercycle length and frequent occurrence of normal outbursts, is much higher than that of other SU UMa-type dwarf novae with similar orbital periods.  The authors are grateful to observers of VSNET Collaboration and VSOLJ observers, who contributed observations, covering long-term variations of ER UMa as long as two years.  This work was partly supported by the Grant-in-Aid ``Initiative for High-Dimensional Data-Driven Science through Deepening of Sparse Modeling'' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan, and Grants-in-Aid for the Global COE Programme ``The Next Generation of Physics, Spun from Universality and Emergence'' from the Ministry of Education, Culture, Sports, Science and Technology of Japan. We are very grateful to the referee of this paper, Osaki, Y.    \\appendix"
    },
    "1402/1402.1907_arXiv.txt": {
        "abstract": "We study the interactions between a protostar and its circumstellar disc under the influence of an external binary companion to determine the evolution of the mutual misalignment between the stellar spin and the disc angular momentum axes.  The gravitational torque on the disc from an inclined binary companion makes the disc precess around the binary axis, while the star-disc interaction torque due to the rotation-induced stellar quadrupole tends to make the stellar spin and the disc angular momentum axes precess around each other. A significant star-disc misalignment angle can be generated from a small initial value as the star-disc system evolves in time (e.g., with decreasing disc mass) such that the two precession frequencies cross each other. This ``secular resonance'' behaviour can be understood in a simple, geometric way from the precession dynamics of spin and disc angular momenta. We derive the general conditions for such resonance crossing to occur, and find that they can be satisfied for very reasonable protostar-disc-binary parameters. The evolution of star-disc inclination is also significantly affected by mass accretion from the disc onto the central star and by magnetic star-disc interaction torques, which can either promote or reduce star-disc misalignment, as well as by the possible damping of disc-binary inclination due to viscous dissipation of disc warps. In general, as long as the initial binary-disc inclination is not too small (greater than a few degrees), a variety of star-disc misalignment angles can be generated within the lifetimes of protoplanetary discs ($\\sim 10$~Myrs). We discuss the implications of our results for the observations of stellar spin orientations in binaries, for the alignments/misalignments of protostellar discs and debris discs relative to their host stars, and for the observed stellar obliquities in exoplanetary systems. In particular, if hot Jupiters are produced by the secular Lidov-Kozai effect induced by an external stellar companion present in the protostellar phase, then it is likely that ``primordial'' star-disc misalignments are already generated by the star-disc-binary interactions.  Even for systems where the Kozai effect is suppressed, misaligned planets and hot Jupiters may still be produced during the protoplanetray disc phase. ",
        "introduction": "In the last few years, many exoplanetary systems containing hot Jupiters have been be found to have high stellar obliquities, i.e., large misalignment angles between the spin axis of the host star and the planetary orbital angular momentum axis (e.g., Hebrard et al.~2008; Winn et al.~2009,2010; Triaud et al.~2010; Moutou et al.~2011; Albrecht et al.~2012). This indicates that a large population of hot Jupiters are formed through high-eccentricity channels, either via dynamical planet-planet scatterings (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Chatterjee et al.~2008; Juric \\& Tremaine 2008) or/and secular interactions between multiple planets, or Lidov-Kozai effect induced by a distant companion (e.g., Holman et al.~1997; Wu \\& Murray 2003; Fabrycky \\& Tremaine 2007; Wu et al.~2007; Nagasawa et al.~2008; Naoz et al. 2011,2102; Wu \\& Lithwick 2011; Katz et al.~2011). Alternatively, Rogers et al.~(2012) suggested that the observed stellar obliquities may result from the torques of varying directions deposited on the radiative stellar envelope by internally generated gravity waves.  The observed stellar obliquities may also contain contributions from (or be affected by) possible ``primordial'' misalignment between the the stellar spin and the protoplanetary disc. Several lines of evidence suggest that such primordial misalignment is possible:  (i) In the Solar system, all major planets lie in the same plane to within $2^\\circ$, but this plane is inclined with respect to the solar rotational equator by $7^\\circ$.  (ii) Solar-type main-sequence binaries with separations $\\go 40$~AU often have rotation axis misaligned relative to the orbital angular momentum (Hale 1994).  (iii) Some binary young stellar objects (YSOs) are observed to have circumstellar discs that are misaligned with the binary orbital plane (e.g., Stapelfeldt et al.~1998,2003; Neuhauser et al.~2009). Also, several unresolved YSOs or pre-main sequence binaries have jet axes along different directions compared to the binary axes, suggesting misaligned discs (e.g., Davis, Mundt \\& Eisl\\\"offel 1994; Roccatagliata et al.~2011).  (iv) Recently, stellar obliquities in several multiplanet systems have been measured or constrained. While for some systems (Kepler-30, Sanchis-Ojeda et al.~2012; KOI-94 and Kepler-25, Albrecht et al.~2013, Hirano et al.~2012; Kepler-50 and Kepler-65, Chaplin et al.~2013), the stellar obliquities are consistent with zero (with typical error bars of order 10$^\\circ$), highly misaligned systems have also been found.  In particular, Huber et al.~(2013) used asteroseismology to measure a large stellar obliquity ($40^\\circ$-$55^\\circ$) for Kepler-56, a red giant star hosting two transiting coplanar planets with orbital period of 10.5 and 21 days. A companion in a wide orbit (as suggested by radial velocity data) may be responsible for the observed stellar obliquity; but it is possible that the companion was already present in the protoplanetary disc phase and produced a primordial misalignment.  Also, there exists tentative evidence that the Kepler-9 system, containing three transiting, roughly coplanar planets (Holman et al.~2010), has a misaligned stellar spin (about 45$^\\circ$) with respect to the planetary orbit (Walkowicz \\& Basri 2013). Hirano et al.~(2014) presented several other candidate misaligned multiplanet systems.   In contrast to the above, the current observational situation with debris disc systems seems different. In a number of debris disc systems, the stellar spin inclination (relative to the line of sight) $i_\\star$ is similar to the debris disc inclination $i_{\\rm disc}$, typically to within 10$^\\circ$ (Watson et al.~2011; Greaves et al.~2013; Kennedy et al.~2013). While $i_\\star\\approx i_{\\rm disc}$ does not necessarily imply alignment between the spin axis and the disc axis, statistically these observations do suggest that there is no significant star-disc misalignment in these systems.  There are several proposed mechanisms for tilting a star relative to its protoplanetary disc to produce primordial misalignment. Bate et al.~(2010) suggested that in the chaotic star formation scenario, the accreting gas assembled onto a protoplanetary disc may have a varying direction of angular momentum, and thus the stellar spin direction can be different from the disc orientation at the later phase of accretion. Misalignment may also be produced when a protostellar system encounters another disc/envelope system with a different direction of rotation (Thies et al.~2011). Lai, Foucart \\& Lin (2011) showed that magnetic star-disc interaction produces a misalignment torque between the stellar spin and the disc axis (see also Lai 1999). However, this torque is of the same order of magnitude as the accretion torque, which tends to align the spin with the disc. Given the intrinsic complexity of the magnetic accretion physics, the net effect of the accretion and magnetic torques cannot be definitively determined at present: All one can say is that small or modest spin-disc misalignments may be produced in some systems; perhaps this is responsible for the $7^\\circ$ anomaly in the Solar system. To generate larger misalignments or retrograde discs require additional external perturbations (such as varying orientations of outer discs; see Foucart \\& Lai 2011).  Recently, Batygin (2012) suggested that the precession of a protoplanetary disc driven by the gravitational torque from an inclined binary companion can lead to misalignment between the stellar spin and the disc axis. This scenario is appealing since the same inclined binary companion is needed to make misaligned hot Jupiters through the secular Lidov-Kozai effect. Batygin \\& Adams (2013) further considered the coupling between the rotating central star and the precessing disc, and showed that significant misaligment between star and disc can be produced.  Since a large fraction of protostars are in binaries, the scenarios examined by Batygin (2012) and Batygin \\& Adams (2013) are of great interest and the mechanism of producing misalignment may be very robust.  On the other hand, as noted above, the observational situation concerning spin-disc misalignment is confusing at present, with some protoplanetary discs showing misalignment, while many debris disc systems showing no significant misalignment.  In this paper we study the interactions between a protoplanetary disc, its host star and an external binary companion on an inclined orbit in order to understand the possibility and conditions of generating large spin-disc misalignments. Our approach to this problem is different from that of Batygin (2012) and Batygin \\& Adams (2013), and we also examine several additional physical effects not considered before.  In particular, we model the star-disc-binary gravitational interactions directly using angular momentum equations (Sections 2-3).  This allows us to describe the rotational dynamics of the system clearly in a physical way and see how different behaviours may arise under different conditions (i.e., different parameters of the system, such as disc size, rotation of the central star and the binary separation). We elucidate the significance of ``secular resonance'' in a geometric way (see Fig.~2 below) -- such resonance plays an important role in producing large spin-disc misalignments.  In Section 4, we study the effects of mass accretion and magnetic interaction between the star and the disc, including the possibility of magnetically driven spin-disc misalignment. We also briefly examine the effect of disc inclination damping (relative to the binary) associated with viscous dissipation of disc warping (Section 5). Overall, we find that all these effects can influence the final star-disc inclination significantly, and a variety of misalignment angles can be produced when a protostar disc is perturbed by a binary companion. In Section 6, we summarize the key physical effects and results (e.g., the conditions for resonance crossing), and discuss the implications of our findings for stellar binaries, protoplanetary/debris discs and the formation of misaligned planets. ",
        "conclusions": "In this paper we have examined the idea (Batygin 2012; Batygin \\& Adams 2013) that the tidal torque on a circumstellar disc from an external binary companion can lead to the generation of misalignment between the disc angular momentum axis and the stellar spin axis. Such ``primordial'' spin-disc misalignment would contribute to the spin-orbit misalignments observed in many exoplanetary systems containing hot Jupiters. We model the star-disc-binary interactions directly using angular momentum equations, which allow us to describe the secular dynamics of the star-disc-binary system in a simple way and derive the conditions (e.g., disc size, binary separation and stellar rotation rate) under which significant spin-disc misalignments can be produced.  \\subsection{Key Physical Effects and Results}  In general, the tidal torque from the external binary companion makes the circumstellar disc precess around the binary angular momentum axis (at the rate $\\Omega_\\rmpd$) -- the disc behaves approximately as a rigid body because the different regions of the disc are coupled by internal waves, viscous stresses or self-gravity.  When the disc is misaligned with the rotation axis of the central star, it tends to drive the stellar spin into precession around the disc axis (at the rate $\\Omega_\\rmps$) due to the rotation-induced oblateness of the star (see Fig.~1). If $|\\Omega_\\rmps/\\Omega_\\rmpd|\\gg 1$, the stellar spin axis $\\hatS$ will follow the disc axis $\\hatL_\\rmd$ adiabatically, so that the spin-disc misalignment angle $\\thesd$ remains approximately constant.  However, as the star-disc system evolves in time (e.g., due to decreasing disc mass), $\\Omega_\\rmps$ decreases [see eq.~(\\ref{eq:Omegaps})].  When the ratio $|\\Omega_\\rmps/\\Omega_\\rmpd|$ crosses unity, significant spin-disc misalignment can be generated. The importance of this secular resonance can be understood in a geometric way (see Fig.~2 and Section 3.4). When $|\\Omega_\\rmps/\\Omega_\\rmpd|\\ll 1$, the stellar spin effectively precesses around the binary axis $\\hatL_\\rmb$ with constant spin-binary inclination angle $\\thesb$.  Thus, if the stellar spin and the disc axis are aligned initially, significant spin-disc misalignment can be produced only if the system experiences resonance crossing during its evolution \\footnote{If the system is in the ``non-adiabatic'' regime $|\\Omega_\\rmps/\\Omega_\\rmpd|\\ll 1$ at all times, then then $\\thesd$ will simply oscillate between $0^\\circ$ and $2\\thedb$, so can periodically attain a large value if $\\thedb$ is large (see Fig.~1). However, this is unlike for typical astrophysical star-disc-binary parameters of interest. See equation (\\ref{eq:con1}).}. In the simple disc evolution model considered in this paper (see Sections 2.1 and 3.3), in order to satisfy $|\\Omega_\\rmps/\\Omega_\\rmpd|\\go 1$ at $t=0$ (with the initial disc mass $M_{\\rm di}$), we require [see eq.~(\\ref{eq:ratio})] \\be {a_\\rmb\\over \\rout}\\go 2.8\\left({M_\\rmb\\over M_\\star}\\right)^{\\!\\!1/3} \\!\\!\\left({\\rout\\over 50\\,{\\rm AU}}\\right)^{\\!\\!-1/6} \\!\\left(\\!{\\bar\\Omega_\\star\\over 0.1}\\!\\right)^{\\!\\!-7/9} \\!\\left(\\!{M_{\\rm di}\\over 0.1M_\\star}\\!\\right)^{\\!\\!-1/3}, \\label{eq:con1}\\ee where $a_\\rmb$ is the binary separation, $\\rout$ is the outer disc radius, and we have ignored the non-essential factors on the second line of equation~(\\ref{eq:ratio}), and have assumed that $\\brin$ and $\\bar\\Omega_\\star$ are related by $f_\\star=0.8$ [see eq.~(\\ref{eq:fastness})] and thus $(\\bar\\Omega_\\star/0.1)\\brin^{-2}=(\\bar\\Omega_\\star/0.1)^{7/3}$. On the other hand, to satisfy $|\\Omega_\\rmps/\\Omega_\\rmpd|\\lo 1$ at $t=10$~Myrs (with the ``final'' disc mass $M_{\\rm df}$), we require \\be {a_\\rmb\\over \\rout}\\lo 7.6\\left({M_\\rmb\\over M_\\star}\\right)^{\\!\\!1/3} \\!\\!\\left({\\rout\\over 50\\,{\\rm AU}}\\right)^{\\!\\!-1/6} \\!\\left(\\!{\\bar\\Omega_\\star\\over 0.1}\\!\\right)^{\\!\\!-7/9} \\!\\left(\\!{M_{\\rm df}\\over 0.005M_\\star}\\!\\right)^{\\!\\!-1/3}. \\label{eq:con2}\\ee Thus, in order for produce significant spin-disc misalignment through resonance crossing during the 10~Myrs of the disc lifetime (more precisely when the disc mass changes from $0.1M_\\star$ to $0.005M_\\star$), the conditions (\\ref{eq:con1})-(\\ref{eq:con2}) must be satisfied. Interestingly, most circumstellar disc systems with observed/imaged binary companions have $a_\\rmb/\\rout$ in the range between a few to 10 (e.g., Stapelfeldt et al.~1998,2003; Neuhauser et al.~2009; Karl Stapelfeldt, private communication 2013) -- such a range is also largely expected from theoretical considerations of the tidal truncation of discs in binaries (e.g., Artymowicz \\& Lubow 1994).  So these conditions can be met in general, although may not always. We note that even when these conditions are not satisfied, $\\thesd$ can still experience modest growth (see, e.g., Fig.~6 and the lower panel of Fig.~7).  Another important aspect of our paper is to incorporate the effects of accretion and magnetic torques in the evolution of the stellar spin direction (Section 4). In general, the magnetic torques between a protostar and its circumstellar disc can induce spin-disc misalignment and contribute to the star-disc mutual precession (Lai et al.~2011), while accretion tends to damp the misalignment.  Unfortunately, given the intrinsic complexity of the magnetosphere-disc interaction physics, the net effect accretion and magnetic torques cannot be determined with certainty at present. Nevertheless, we show that these torques in general can influence the evolution of the stellar spin axis in a significant way (see Figs.~8,9,11).  Finally we briefly explore the effect of disc inclination damping (relative to the binary) associated with viscous dissipation of disc warps (Section 5).  Only with extreme damping timescales (less than 1 Myrs) can the spin-disc misalignments be significantly affected.  By comparing with the observations of spin-orbit misalignments in wide binaries and spin-disc misalignments in protostellar systems, it may be possible to constrain the physics of disc warp dissipations.    \\subsection{Implications}  The general implication of our work is that in the presence of a binary companion, stellar spin-disc misalignments can be generated within the typical lifetimes of protoplanetary discs under a wide range of conditions, as long as the binary axis is somewhat misaligned with the disc axis (by more than a few degrees). There is ample evidence of such disc-binary misalignments from observations of protostellar jets or direct imagings of protostellar discs (see Section 1). In fact, for very reasonable binary/disc parameters, it has proven difficult to avoid the production of significant spin-disc misalignments from small initial values.  Although challenging, it would be of great interest to measure or constrain the spin-disc inclination angles of young protostellar disc systems (rather than debris disc systems; see below).  Our general conclusion is consistent with the observations of spin-orbit misalignments in wide ($a_\\rmb \\go 40$~AU) binaries (Hale 1994). Since there is significant evolution of the spin-disc misalignment angle during the lifetime of the disc, our work suggests that the observed spin-orbit misalignments do not necessarily translate into a large difference in the angular momentum directions of separate molecular cloud cores that form the binary components.  Concerning giant planets, our work suggests that spin-disc misalignments generated by star-disc-binary interactions can make a significant contribution to the observed spin-orbit misalignments in hot Jupiter systems. Such misalignments may already be present prior to any subsequent (after disc dispersal) dynamical processes. Indeed, the same kind of inclined binaries are invoked in the Lidov-Kozai mechanism for producing hot Jupiters (see references in Section 1). Moreover, in order for the Lidov-Kozai migration to operate on a planet at the semi-major axis $a_p$, the Kozai oscillation period ($\\simeq 2\\pi/\\dot\\omega_{\\rm Kozai}$) \\be P_{\\rm Kozai}\\simeq 10^6\\left({a_\\rmb\\over 500\\,{\\rm AU}}\\right)^3 \\left({a_p\\over 5\\,{\\rm AU}}\\right)^{\\!\\!-3}\\,{\\rm yrs} \\ee must be shorter than the apsidal precession period of the planetary orbit due to general relativity (e.g., Holman et al.~1997). That is, the ratio \\be {\\dot\\omega_{\\rm GR}\\over \\dot\\omega_{\\rm Kozai}}\\simeq 6\\times 10^{-3} \\left(\\!{M_\\star\\over M_\\rmb}\\!\\right)\\left({a_\\rmb\\over 500\\,{\\rm AU}}\\right)^{\\!3} \\left({a_p\\over 5\\,{\\rm AU}}\\right)^{\\!\\!-4} \\ee must be less than unity. Thus, if the planet was formed at $a_p\\lo 1.4$~AU or had migrated (due to disc-driven migration) to such a distance by the end of the protoplanetary disc phase, the binary companion (assuming $a_\\rmb =500$~AU) would not be able to induce any Kozai oscillation. Nevertheless, such ``failed-Kozai'' systems may still have significant spin-orbit misalignments because of the primordial spin-disc misalignments studied in this paper.  Of course, whether such systems can have any significant eccentricities or can become hot Jupiters is a different question. In this regard, it is of interest to note that there appears to be a significant lack of super-eccentric proto-hot Jupiters in the Kepler sample (Dawson et al.~2012), suggesting that many hot Jupiters are not produced by high-eccentricity migration mechanism (see Socrates et al.~2012; Dawson \\& Murray-Clay 2013).  To predict the planet's orbital inclination at the end of the protoplanetary disc phase, it will be necessary to keep track of the evolution of the planet's angular momentum (both magnitude and direction) as the disc evolves and precesses, together with all the other effects considered in this paper. Planet-disc interaction will be an important part of the story (e.g., Baruteau et al.~2013; see also Marzari \\& Nelson 2009; Terquem 2013; Xiang-Gruess \\& Papaloizou 2013).  Presumably, when the disc is sufficiantly massive as the planet forms and evolves, the angular momentum of the planet will be strongly coupled to that of the disc. If the planet migrates to the close vicinity of the star, the mutual torque between the star and the planet must also be considered.  We plan to address some of these issues in the future.  Concerning debris disc systems, current observations do not reveal any evidence of significant spin-disc misalignments (Watson et al.~2011; Greaves et al.~2013; Kennedy et al.~2013). This suggests that the observed/measured systems, so far, do not have binary companions at all or companions with appropriate separations. Finding binary companions in debris disc systems and measure/constrain the relative orientations of spin, disc and binary axes would be of great interest. On the theoretical side, it would be necessary to extend our study to much later times (beyond 10~Myrs).  This would require more reliable models for the evolution of discs from the protoplanetary phase to the debris disc phase."
    },
    "1402/1402.6699_arXiv.txt": {
        "abstract": "A planet's emission spectrum contains information about atmospheric composition and structure. We compare the Bayesian Information Criterion (BIC) of blackbody fits and idealized spectral retrieval fits for the 44 planets with published eclipse measurements in multiple thermal wavebands, mostly obtained with the Spitzer Space Telescope. The evidence for spectral features depends on eclipse depth uncertainties. \\emph{Spitzer} has proven capable of eclipse precisions better than $10^{-4}$ when multiple eclipses are analyzed simultaneously, but this feat has only been performed four times.  It is harder to self-calibrate photometry when a single occultation is reduced and analyzed in isolation; we find that such measurements have not passed the test of repeatability.  Single-eclipse measurements either have an uncertainty floor of $5\\times10^{-4}$, or their uncertainties have been underestimated by a factor of 3. If one adopts these empirical uncertainties for single-eclipse measurements, then the evidence for molecular features all but disappears: blackbodies have better BIC than spectral retrieval for all planets, save HD~189733b, and the few planets poorly fit by blackbodies are also poorly fit by self-consistent radiative transfer models. This suggests that the features in extant broadband emission spectra are due to astrophysical and instrumental noise rather than molecular bands.  Claims of stratospheric inversions, disequilibrium chemistry, and high C/O ratios based solely on photometry are premature. We recommend that observers be cautious of error estimates from self-calibration of small data sets, and that modelers compare the evidence for spectral models to that of simpler models such as blackbodies. ",
        "introduction": "An exoplanet on an edge-on orbit periodically passes behind its host star.  The decrement in thermal flux that occurs during such an eclipse is a measure of the dayside brightness temperature of the planet.  The brightness temperature of a planet varies with wavelength, primarily because of the atmosphere's wavelength-dependent opacity and vertical temperature profile \\citep[e.g.,][]{Deming_2005, Seager_2005, Barman_2005, Burrows_2007, Burrows_2008, Fortney_2008, Knutson_2008, Desert_2009}. If different wavelengths probe the same atmospheric layer (e.g., a cloud deck) then the planet will appear to have a blackbody spectrum. In the absence of clouds, a planet may still have a blackbody spectrum if the atmospheric layers probed are isothermal.  Indeed, the emission spectra of some planets are reported to be featureless: e.g., TrES-2 \\citep{ODonovan_2010}, TrES-3 \\citep{Fressin_2010}, WASP-18b \\citep{Nymeyer_2011}, and WASP-12b \\citep{Crossfield_2012}.  In principle, the detection of molecular bands in the infrared emission spectrum of a planet enables the retrieval of greenhouse gas abundances and the vertical temperature profile of the planet \\citep[e.g.,][]{Madhusudhan_2009b, Madhusudhan_2011, Lee_2012, Line_2012}.  Spectral resolution is critical to such retrieval exercises because a high-resolution emission spectrum is more likely to deviate significantly from a blackbody, and renders the retrieval problem well-constrained.  This bodes well for current and future efforts to perform \\emph{bona fide} emission spectroscopy.  So far, however, the vast majority of exoplanet emission measurements have been broadband eclipse photometry.  A typical retrieval model uses a dozen parameters to describe the atmospheric composition and vertical temperature profile, while a typical hot Jupiter has only been observed in 2--4 thermal broadbands.  Even for the few planets with 6 or 7 thermal measurements, the photometric retrieval problem is under-constrained.  A widely noted consequence of the parameter--data mismatch is that exact atmospheric properties cannot be uniquely determined, making color-color and color-magnitude diagrams more realistic approaches to atmospheric classification \\citep{Baskin_2013, Beatty_2014, Triaud_2014}.  The less-discussed aspects of under-constrained retrieval are that (1) there is no way to reject erroneous measurements, and (2) the estimated uncertainties on eclipse depths directly affect the uncertainties on atmospheric parameters.  This is in stark contrast to over-constrained problems such as fitting an occultation model to time-series data, for which it is customary to perform outlier rejection (e.g., $\\sigma$-clipping), and for which the photometric uncertainties are estimated in the process of fitting a model to the data, rather than trusting the output of aperture photometry routines.  Nonetheless, many exoplanet discoveries have been based on broadband emission spectra: a temperature inversion in the atmosphere of HD~209458b was inferred from 4 broadband eclipse depths \\citep{Knutson_2008}, disequilibrium chemistry was invoked to explain the 6-band emission spectrum of GJ~436b \\citep{Stevenson_2010}, and high atmospheric C/O was discovered based on 7 broadband eclipses of WASP-12b \\citep{Madhusudhan_2011}.  These successes have led to classifying planets based on temperature inversions \\citep[using 2 broadbands per planet;][]{Knutson_2010} and C/O ratio \\citep[using $\\ge 4$ bands per planet;][]{Madhusudhan_2012}.  Temperature inversions and non-solar chemistry have since been disputed for each of the exemplar planets due to the re-reduction of existing data \\citep{Beaulieu_2011} acquisition of new data at the same wavelength \\citep{Cowan_2012, Zellem_2014} or acquisition of new data at different wavelengths \\citep{Crossfield_2012}.  Such challenges are not unique to eclipse \\emph{photometry}: the featureless day-side emission spectrum of HD~189733b \\citep{Grillmair_2007} exhibited an absorption feature at a later epoch \\citep{Grillmair_2008},\\footnote{Although this was interpreted as evidence of planetary variability, that hypothesis is inconsistent with the more extensive monitoring campaign of \\cite{Agol_2010}.} and line emission from the dayside of HD~189733b \\citep{Swain_2010} has been disputed by \\cite{Mandell_2011}.  Nor are issues of repeatability limited to superior conjunction: the first thermal phase measurements of an exoplanet \\citep{Harrington_2006} were later found to be off by $80^\\circ$ in phase and more than a factor of 2 in amplitude \\citep{Crossfield_2010}; the first half of the thermal phase measurements of \\cite{Knutson_2007b} were later found to be corrupted by detector systematics \\citep{Agol_2010}.  The situation is similar for transit spectroscopy, where initial claims of molecular absorption bands \\citep{Tinetti_2007, Swain_2008, Tinetti_2010} were disputed on the basis of data reduction, error estimation, and astrophysical variability \\citep{Ehrenreich_2007, Desert_2009, Gibson_2011, Desert_2011c, Crouzet_2012}.  Indeed, \\cite{Burrows_2014} offers a sobering review of the exoplanet atmospheric characterization field, speculating that many of the extraordinary claims of the past decade may be overturned by better data.  In this article we attempt to reconcile Burrows' pessimistic view with the growing body of papers making statements about planetary atmospheres based on a handful of eclipse measurements.  Instead of focusing on a single planet, we perform a holistic analysis of all transiting planets with multiple eclipse measurements.  We consider only broadband measurements (for which it is easy to quantify the number of independent observational constraints) of eclipse depths (which are unaffected by star spots). Our approach is to compare the goodness-of-fit and evidence for three classes of models: blackbodies, self-consistent radiative transfer, and spectral retrieval.  Since the disputes over atmospheric properties have often revolved around the reliability of eclipse depths, we empirically estimate the accuracy of broadband eclipse measurements. Notably, the dominant ``signal'' in space-based eclipse photometry is usually the detector sensitivity, which must be modeled using the very same observations of the science target.  Future observations of transiting planets with the James Webb Space Telescope are likely to resolve many of the current scientific disputes about the nature of hot Jupiter atmospheres. Attempts to push the observatory to smaller and cooler planets, however, will still rely on self-calibration; error estimation and repeatability will therefore remain critical issues. ",
        "conclusions": "The retrieval of atmospheric structure and composition from disk-integrated broadband photometry hinges on planets not emitting like blackbodies. We have considered the 44 short-period planets with emission measurements in multiple broadbands.  If published uncertainties are taken at face value, then seven of these planets have broadband spectra that favor spectral retrieval over blackbody fits, according to the Bayesian Information Criterion---CoRoT-2b benefits the most from the additional model parameters.  In order to perform under-constrained spectral retrieval, however, it is critical to know the actual uncertainty on eclipse measurements. \\emph{Spitzer} is capable of exquisite photometry ($<10^{-4}$), but single eclipses acquired, reduced and analyzed in isolation have only been repeatable at the $1\\sigma$ level of $5\\times10^{-4}$ (or single-eclipse uncertainties have been under-estimated by a factor of 3). If one adopts such empirical uncertainties for single-eclipse measurements, then blackbody fits are preferable over spectral retrieval for all planets, with the possible exception of HD~189733b.  We conclude that statements about atmospheric composition based solely on broadband emission measurements are premature.  If one adopts empirical estimate of single-eclipse accuracy, then HD~209458b and GJ~436b are well fit by blackbodies, and WASP-12b is not so poorly fit as to favor spectral retrieval. This resonates with the cautionary review of \\cite{Burrows_2014}. Temperature inversions and odd compositions were inferred for short period planets based on broadband emission spectra \\citep{Knutson_2008, Knutson_2010, Stevenson_2010, Madhusudhan_2011, Madhusudhan_2012}. Our results call these claims into question. Undoubtedly, many planets have stratospheric inversions and non-solar chemistry, but there is no robust evidence for this in the current photometry of short-period planets."
    },
    "1402/1402.3769_arXiv.txt": {
        "abstract": "We present a catalog of 1921 spectroscopic redshifts measured in the fields of the massive galaxy clusters MACSJ0416.1--2403 ($z=0.397$), MACSJ0717.5+3745 ($z=0.546$), and MACSJ1149.5+2223 ($z=0.544$), i.e., three of the four clusters selected by STScI as the targets of the  Frontier Fields (FF) initiative for studies of the distant Universe via gravitational lensing. Compiled in the course of the MACS project (Massive Cluster Survey) that detected the FF clusters, this catalog is provided to the community for three purposes: (1) to allow the identification of cluster members for studies of the galaxy population of these extreme systems, (2) to facilitate the removal of unlensed galaxies and thus reduce shear dilution in weak-lensing analyses, and (3) to improve the calibration of photometric redshifts based on both ground- and spacebased observations of the FF clusters. ",
        "introduction": "Recognizing the importance of massive galaxy clusters as gravitational telescopes for studies of the distant universe as originally proposed by the Clusters as Telescopes team (CATs\\footnote{Led by Jean-Paul Kneib, the CATs team consists of cluster and lensing experts from the US and Europe; CATs was one of six teams selected by STScI in the spring of 2013 to calibrate the FF cluster lenses.}), Space Telescope Science Institute (STScI) recently announced the Frontier Fields (FF) project, a commitment of 560 orbits of \\textit{Hubble Space Telescope (HST)} time in Cycles 21 and 22 to deep observations of four clusters of galaxies at intermediate redshift ($0.3<z<0.6$). Each target will be observed with  both the Advanced Camera for Surveys (ACS) and the Wide Field Camera 3 (WFC3) for 20 orbits in each of the F435W, F606W, F814W, F105W, F125, F140, and F160W filters, reaching a limiting magnitude of approximately 29 (AB system) across the full optical and near-infrared (NIR) window. Both ACS and WFC3 will be used simultaneously, such that for each cluster target images of equal depth are accumulated of the cluster core and of a flanking parallel field about 6 arcmin away. Extending the NIR spectral coverage further,  1000 hours of Spitzer Space Telescope time are being devoted to observations of the same four fields at 3.6 and 4.5$\\mu$m.  The FF project represents the  largest investment of  \\textit{HST} time for deep observations of galaxy clusters in the history of  \\textit{HST}. Its main science goal is the discovery of gravitationally lensed background galaxies at $z=5-10$ for both in-depth studies of bright individual objects and statistical investigations into the properties of galaxies at magnitudes and distances inaccessible to observation without gravitational amplification. As a bonus, the FF observations will also yield the deepest images of the cores of massive clusters ever collected, thereby facilitating exquisitely detailed characterizations of the cluster lenses and their galaxy content.  The four clusters selected for the FF initiative for Cycle 21 and 22 are A2744 ($z{=}0.308$), MACSJ0416.1--2403 ($z{=}0.397$), MACSJ0717.5+3745 ($z{=}0.546$), and MACSJ1149.5+2223 ($z{=}0.544$) \\citep{1958ApJS....3..211A,2003MNRAS.339..913E,2007ApJ...661L..33E,2010MNRAS.407...83E,2012MNRAS.420.2120M}. Subject to a review of the observations collected up to then, two additional clusters (AS1063 and A370) will be added to the FF sample and observed in Cycle 23. All FF clusters are exceptionally X-ray luminous systems ($L_{\\rm X}>2\\times 10^{45}$ erg s$^{-1}$, bolometric) and feature pronounced substructure, i.e., they combine the two characteristics that select the most powerful cluster lenses \\citep{2010MNRAS.406.1318H,2012A&A...544A..71L,2013ApJ...762L..30Z}. Additional information on the FF project can be found on \\url{http://www.stsci.edu/hst/campaigns/frontier-fields/}. ",
        "conclusions": ""
    },
    "1402/1402.1593_arXiv.txt": {
        "abstract": "We have made an extensive search for grouping  amongst  the near Earth asteroids (NEAs). We used two D- functions and  rigorous cluster analysis approach. We have found several new groups (associations) among the NEAs:  the objects moving on similar orbits with small minimum orbital intersection distances (MOID) with the Earth trajectory. Reliability of some of these groups is quite high. \\vspace{-6pt} ",
        "introduction": "% The existence of the main belt asteroid families is beyond the doubts. However existence of groups  among the NEAs,  with members of common origin, suggested by similarity between their orbital  parameters \\citep{Drummond1991, Obrubov1991, Fu2005, Jopek2011, Schunova2012} is not yet generally accepted.  Many groups found by these authors might be attributed to chance alignments.  In this study we have searched for associations amongst $9004$ NEA's osculating orbits \\citep{Neodys2012} by two strict cluster analysis methods, similar to that used in \\cite{Jopek2011}. Two D-functions were used: $D_{SH}$ introduced in \\cite{SouthworthHawkins1963}  and $D_H$ introduced in \\cite{Jopek1993}. The orbital similarity thresholds corresponded to $1$\\% probability that given group was found by chance alignment. The clusters were detected by a single neighbour linking technique. \\vspace{-12pt} ",
        "conclusions": "Using two searching methods we found $20$ groups of ten or more members: $13$ groups with the $D_{SH}$ function and $9$ groups using $D_H$ function.  Among 20 formally separated groups we have selected 10 associations.  All but two has been found in both searches, however with different amount of members. Associations No 1 and 2 (SH -179 and SH-380)  were detected only with $D_{SH}$ function. Of course when the values of the similarity thresholds  has been slightly increased, these groups were found also with $D_H$ function. Associations No 3-7 were identified with both functions with considerable amount of members in common. In case of associations 8,9 and 10 the results shown to be more complicated: several overlapping subgroups were detected for which the common members were less numerous. In Table \\ref{meanelements} we see that for all groups the mean orbital inclinations are smaller than $10$ degrees. It is not unexpected, we should recall that all NEAs but two have $i<75$ degrees, and for $\\sim 4500$ orbits the inclinations are smaller than $10$ degrees.  \\begin{table} \\caption{The list of $10$ NEAs associations of $10$ or more members found in this study. The codes of the associations include the distance function tag and the ordinal number of the asteroid from the NEODyS list which was identified as a first group member. The name of this asteroid is also given. $N$ is the amount of members in the group. In brackets the amount of the common members identified by both functions are given. Within each group the orbital elements were averaged by the method described in \\cite{Jopek2006}. } \\centering \\scriptsize \\begin{tabular}{|r|l|l|r|r|r|r|r|r|r|} \\hline \\multicolumn{1}{|c}{No}              & \\multicolumn{1}{|c} {NEA }             & \\multicolumn{1}{|c}{Code }           & \\multicolumn{1}{|c}{N}               & \\multicolumn{1}{|c}{ a [AU] }         & \\multicolumn{1}{|c }{ q [AU]}         & \\multicolumn{1}{|c| }{e }             & \\multicolumn{1}{|c| }{i [deg] }             & \\multicolumn{1}{|c| }{$\\omega$ [deg] }             & \\multicolumn{1}{|c| }{$\\Omega$ [deg] }            \\\\ \\hline 1 & '13553' Masaakikiyama &SH-179  &    53       &  2.133 &   1.154 &   0.459 &     5.7 &   157.3 &   143.5 \\\\ 2 & '89136' 2001US16      &SH-380  &    24       &  1.375 &   1.025 &   0.254 &     1.0 &    16.0 &   227.3 \\\\ 3 & '2368' Beltrovata     & SH-35  &    77(74)   &  2.128 &   1.221 &   0.426 &     2.7 &    32.2 &   301.6 \\\\ &                       &  H-35  &   101       &  2.137 &   1.222 &   0.428 &     3.0 &    31.2 &   305.6 \\\\ 4 & '4660' Nereus         & SH-77  &    63(40)   &  1.864 &    .936 &   0.498 &     1.2 &    37.2 &    82.1 \\\\ &                       & H-77   &    63       &  1.793 &   0.959 &   0.465 &     1.0 &    46.0 &    70.9 \\\\ 5 & '8014' 1990MF         & H-143  &    69       &  2.019 &   1.051 &   0.479 &     0.1 &   100.3 &   232.1 \\\\ & '10860' 1995LE        & SH-167 &    65(61)   &  2.135 &   1.111 &   0.480 &     1.3 &   160.9 &   167.4 \\\\ 6 & '11054' 1991FA        & H-168  &    68       &  1.513 &   0.984 &   0.350 &     0.9 &    84.7 &   349.1 \\\\ & '190491' 2000FJ10     &SH-866  &    25(23)   &  1.427 &   1.014 &   0.289 &     2.6 &   184.9 &   249.8 \\\\ 7 & '36017' 1999ND43      &SH-255  &    39(21)   &  1.425 &   1.012 &   0.290 &     2.6 &   148.5 &   227.1 \\\\ & '256004' 2006UP       & H-1050 &    35       &  1.421 &   1.013 &   0.287 &     2.8 &   157.4 &   217.0 \\\\ 8 & '54509' YORP          & H-295  &    78       &  1.045 &   1.031 &   0.013 &     1.0 &   321.6 &   224.2 \\\\ & '65717' 1993BX3       & SH-306 &    89       &  1.162 &   0.957 &   0.176 &     0.6 &   235.4 &   210.2 \\\\ & '209215' 2003WP25     & SH-910 &    23       &  1.082 &   0.926 &   0.144 &     0.8 &    95.9 &   188.8 \\\\ 9 & '19356' 1997GH3       &  H-202 &    73       &  1.815 &   1.007 &   0.445 &     0.9 &    14.5 &   145.1 \\\\ & '27002' 1998DV9       & SH-236 &    27(21)   &  1.974 &   1.015 &   0.486 &     3.3 &    10.4 &   130.3 \\\\ & '1994GV'              & SH-1385&    85       &  1.676 &   1.015 &   0.394 &     0.1 &   296.2 &   243.3 \\\\ & '2003DW10'            & H-2582 &    16       &  1.587 &   0.998 &   0.371 &     0.2 &   127.5 &    74.8 \\\\ 10 & '136564' 1977VA       & SH-423 &    17(1)    &  2.172 &   1.141 &   0.475 &     4.2 &   204.1 &   199.3 \\\\ & '2061' Anza           &   H-23 &   181       &  2.202 &   1.046 &   0.525 &     0.2 &   201.2 &   182.1 \\\\ & '4015' Wilson-Harrington& SH-60  &   196(134)  &  2.150 &   1.046 &   0.514 &     0.3 &   169.0 &   223.1 \\\\ \\hline \\end{tabular} \\normalsize \\label{meanelements} \\end{table} We have continued our study gradually decreasing thresholds. As was expected, amount of members of the group was decreasing; some groups were splitting and finally disappearing. Associations SH-306, SH-1385, H23 have proved to be the most robust one. On Figure 1, just for an example we have plotted the orbits of the members of association SH380. \\begin{figure} \\centerline{\\includegraphics[width=4.0cm]{roSH_380_24b}} \\footnotesize \\caption{ Association SH-380 plotted on the ecliptic plane. The orbits remarkably resemble a meteoroid stream. Association includes 24 NEAs: 89136, 26430 8,1994CJ1, 2003CC, 2003GA, 2004KG17, 2005HB4, 2005JT1, 2006HX30, 2006KL103, 2007HB15, 2008GR3, 2008LE, 2009DC12, 2009HG, 2009HH21, 2010CE55, 2011GR59, 2011GV9, 2011OK45, 2011OR5, 2011PU1, 2012GP1, 2012KT12.  The Earth circular trajectory is seen inside the association. } \\normalsize \\end{figure}  In this study we have shown, that using similar rigorous method as for the meteoroid orbits, the NEA's associations can be found easily. However, to ensure about their common origin, the long term numerical integration is needed."
    },
    "1402/1402.2842_arXiv.txt": {
        "abstract": "{ In order to study the temperature distribution of the extended gas within the Orion Kleinmann-Low nebula, we have mapped the emission by methyl cyanide (CH$_3$CN) in its $J$=$6_K$--$5_K$, $J$=$12_K$--$11_K$, $J$=$13_K$--$12_K$, and $J$=$14_K$--$13_K$ transitions at an average angular resolution of $\\sim$10 arcsec (22 arcsec for the $6_K$--$5_K$ lines), as part of a new 2D line survey of this region using the IRAM 30m telescope. These fully sampled maps show extended emission from warm gas to the northeast of IRc2 and the distinct kinematic signatures of the hot core and compact ridge source components. We have constructed population diagrams for the four sets of $K$-ladder emission lines at each position in the maps and have derived rotational excitation temperatures and total beam-averaged column densities from the fitted slopes. In addition, we have fitted LVG model spectra to the observations to determine best-fit physical parameters at each map position, yielding the distribution of kinetic temperatures across the region. The resulting temperature maps reveal a region of hot ($T > 350$~K) material surrounding the northeastern edge of the hot core, whereas the column density distribution is more uniform and peaks near the position of IRc2. We attribute this region of hot gas to shock heating caused by the impact of outflowing material from active star formation in the region, as indicated by the presence of broad CH$_3$CN lines. This scenario is consistent with predictions from C-shock chemical models that suggest that gas-phase methyl cyanide survives in the post-shock gas and can be somewhat enhanced due to sputtering of grain mantles in the passing shock front. } ",
        "introduction": "The Orion Kleinmann-Low (KL) nebula is the closest \\citep[414 pc,][]{Menten2007} and most studied high-mass star-forming region in the Galaxy. The prevailing chemistry in this source is particularly complex as a result of the interactions between newly formed protostars, their outflows and their environment. The evaporation of dust mantles and the presence of high gas temperatures produce a wide variety of molecules in the gas phase that are responsible for a spectacularly rich and intense line spectrum \\citep{Blake1987, Charnley1997}. The chemical complexity of Orion KL has been demonstrated by an extensive number of line surveys, performed over a variety of frequency ranges from 72 GHz up to 1.9 THz \\citep[e.g.,][]{Blake1986, Schilke1997, Bergin2010a}. Finally, millimetre and submillimetre aperture synthesis studies have provided the spatial location and extent of many molecular species \\citep[e.g.,][]{Blake1996, Beuther2005, Plambeck2009, Zapata2009}, further indicating the presence of distinct physical/kinematic components that show clear chemical differentiation, with NH- and OH-bearing molecules peaking in different sources (the hot core and compact ridge, respectively).  Near- and mid-IR subarcsecond resolution imaging and (sub)millimetre interferometric observations have identified the main sources of luminosity, heating, and dynamics in the region \\citep[e.g.,][]{Dougados1993, Gezari1998, Shuping2004, DeBuizer2012}. However, the nature of these objects has been the subject of much debate, as has the question of which source(s) are most responsible for the high luminosity produced by the nebula \\citep[$\\sim$10$^5$ $L_{\\sun}$;][]{Menten1995}. The main sources of activity seem to be radio sources \\textit{I} and \\textit{n}, located just to the south and southwest of IRc2. The Becklin-Neugebauer (BN) object 9$\\arcsec$ northwest of the hot core is believed to be internally heated by an embedded massive protostar (or protostars). Proper motion studies have shown that all three sources, \\textit{I}, \\textit{n}, and BN, are receding from a common origin. This has led to the hypothesis that they were expelled from a massive young stellar system that disintegrated $\\sim$500 years ago \\citep{Bally2005, Zapata2009, Niederhofer2012, Nissen2012}, and that this explosive event may be the source of heating for the region in and around the hot core \\citep{Zapata2009, Zapata2011, Moeckel2012}. The spectacular molecular ``fingers'' that trace the high-velocity outflow emanating from KL in a northwest-southeast direction have also been attributed to this explosive event.  Interferometric observations probing the central region of the KL nebula at arcsecond resolution have also revealed another distinct kinematic component, particularly evident in shock tracers such as SiO and SO, that extends in the northeast-southwest direction and has been identified as a low-velocity bipolar outflow originating at source \\textit{I} \\citep[e.g.,][]{Plambeck1982, Greenhill1998, Niederhofer2012}. This, along with the spectral shape of its continuum emission, has therefore lead various authors to argue that source \\textit{I} is likely to be an embedded protostar driving this outflow. Zapata et al. (2009, 2011) have proposed instead that the low-velocity outflow is linked to the explosive event and that the region surrounding IRc2 and the hot core is being externally heated by material ejected from the explosion impacting the stationary gas. The nature of the Orion hot core, BN object, and sources \\textit{I} and \\textit{n}, and the underlying cause of the dynamics and energetics in the KL nebula are thus matters of continued debate.  Reliable measurements of the temperature distribution across the extended gas in the KL nebula can help to better understand the energetics in the region, and therefore the likely driving sources. Methyl cyanide, CH$_3$CN, is an excellent probe of the gas temperature in warm dense environments like Orion KL \\citep{Boucher1980}. Its molecular symmetry gives rise to many transitions closely spaced in frequency that span a large range of energies (up to $E_\\mathrm{up} \\! > \\! 1000$ K). These lines can therefore be observed simultaneously at one frequency setting, thus avoiding flux calibration uncertainties, and are excited by collisions (so-called $K$-ladder transitions) so that their level populations are dictated by the gas temperature. Nitrogen-bearing complex organic molecules such as methyl cyanide are generally assumed to form on grain mantles in the cold dense phase of the interstellar medium, before being released into the gas phase by evaporation in the hot regions surrounding massive protostars. Their spatial distribution and abundances can therefore provide important constraints on chemical models of these sources.  The methyl cyanide emission in Orion KL has been studied extensively, both with single-dish and interferometric observations. Early observations of the millimetre lines of CH$_3$CN \\citep{Loren1981, Loren1984, Sutton1986} revealed the presence of warm gas ($\\sim$275~K) in the hot core and weaker extended emission from cooler gas ($\\sim$100~K) in the quiescent ridge, with derived fractional abundances of 10$^{-11}$ to 10$^{-10}$. More recent interferometric studies \\citep{Wilner1994, Beuther2006, Wang2010} have uncovered the complex and clumpy structure within $\\sim$10$\\arcsec$ of IRc2, yielding kinetic temperatures above 250~K (and even up to 600~K, based on submillimetre transitions) for the hot core and lower values of $\\sim$150--250~K toward the compact ridge. These arcsecond resolution observations suggest much higher CH$_3$CN abundances of 10$^{-8}$ to 10$^{-7}$ in the hot core and compact ridge.  We have obtained fully sampled maps of four sets of methyl cyanide $K$-ladder emission lines across a large region centred on Orion IRc2. These maps cover a larger area with improved sensitivity and (single-dish) spatial resolution than those of previous studies, providing a homogeneous dataset that is well-suited to studying the extended emission in the region surrounding the KL nebula. Furthermore, while these data cannot compete with the angular resolution afforded by interferometric techniques, they do not suffer from the problem of missing flux common to many interferometric observations, which tend to emphasise the clumpiness of the gas by filtering out the smooth extended emission. Recovering this extended emission is crucial when studying the structure on larger scales, which is the focus of this paper. Based on these data, we have analysed the kinematics of this extended warm gas and derived maps of temperature and column density across the region. We describe the observations that were obtained and the subsequent reduction of those data in Sect.~\\ref{Observations}. The maps are presented in Sect.~\\ref{Results:Emission-Maps} and the general kinematic properties of the emitting gas are discussed. We describe the analysis of these maps and the derivation of column densities and rotational temperatures by use of the population diagram method, followed by LVG model fitting in Sect.~\\ref{Analysis}. The resulting temperature distributions are presented in Sect.~\\ref{Results:Temperature-Maps} and their implications for the source structure are discussed in Sect.~\\ref{Discussion}. Our findings are summarised in Sect.~\\ref{Conclusions}. ",
        "conclusions": "We have obtained full spatially sampled maps of methyl cyanide emission in four sets of its $K$-ladder transitions, $J$=$6_K$--$5_K$, $J$=$12_K$--$11_K$, $J$=$13_K$--$12_K$, and $J$=$14_K$--$13_K$, across the Orion KL region using the IRAM 30m telescope. These maps beautifully display the extended emission from this molecule over an area approximately 40 arcsec across, well beyond the confines of the hot core and compact ridge components. By fitting Gaussian profiles to the line shapes of each set of $K$-ladder transitions, the kinematics of the gas can be studied across the region and reveal a sharp velocity gradient and the distinct kinematic features of the hot core and plateau components. The spectroscopic properties of its emission lines make methyl cyanide an ideal diagnostic of the gas temperature in warm ($>$100 K) and dense ($>$10$^{5}$ cm$^{-3}$) environments. Taking advantage of this, we have used population diagram analysis and LVG modelling to infer the physical conditions [$T_\\mathrm{kin}$, $N_\\mathrm{tot}$, $n(\\mathrm{H_2})$] at each position in our maps. The resulting temperature distributions for the emitting region show a distinct ``hot zone'' surrounding the northeastern edge of the hot core, with inferred temperatures of 300--450 K. This feature is seen in all four sets of $K$-ladder emission lines and emerges from both population diagram and LVG model analyses. The location of the hot zone suggests an association with the northeast-southwest bipolar outflow originating in the vicinity of IRc2. Based on the kinematics and derived physical properties of the methyl cyanide emission in this region, we argue that it is associated with hot shocked gas where the outflow impacts the quiescent material in the extended ridge cloud. We have computed a series of shock chemistry models in order to test this scenario. We find that methyl cyanide can indeed remain present in shock fronts, at high temperatures and elevated abundances. Our main findings are as follows:  \\begin{itemize} \\item Methyl cyanide emission is extended across a 40 arcsec region around IRc2, with temperatures of 100--150 K and inferred densities of order $10^{5}$ cm$^{-3}$ in the quiescent gas. \\item The quiescent gas along the molecular ridge shows a global north-south velocity gradient and narrow line widths ($\\lesssim$4~km\\,s$^{-1}$), while the hot core region at the centre of the maps is clearly identifiable in the CH$_3$CN emission via its distinct kinematic signature, characterised by blue-shifted and broader lines. \\item CH$_3$CN emission peaks at the hot core position; column densities derived from population diagrams and LVG modelling both show that methyl cyanide is over an order of magnitude more abundant in the hot core than in the ambient medium. \\item Temperatures between 200 and 350 K have been determined for the methyl cyanide emission in the hot core and its environs, however, higher temperatures ($\\ge$300 K) are seen in a hot zone to the northeast of the hot core. \\item We attribute this feature to shock heating by outflowing gas from active star formation in the vicinity of IRc2. This is supported by the coincidence of this region with emission from shock tracers such as SiO and neutral carbon. \\item Chemical models of C-shocks indicate that methyl cyanide abundances may be enhanced in shock fronts under physical conditions appropriate for this region. \\end{itemize}  Methyl cyanide is a powerful diagnostic probe of the gas temperature in warm dense regions of star formation. \\hbox{Millimetre} and submillimetre interferometers such as the Plateau de Bure Interferometer, the Submillimeter Array, the Combined Array for Research in Millimeter-wave Astronomy, and the Very Large Array are capable of resolving the clumpy structure within Orion KL and can provide detailed information on its properties, while the Atacama Large Millimeter Array is now able to observe CH$_3$CN with sufficient sensitivity to discern the morphology of the weaker extended emission. Such follow-up observations will be vital to confirm the shock heating scenario we propose here. Combined with legacy datasets such as those presented in this paper, future observations of methyl cyanide promise to unveil the detailed temperature structure in the very hearts of these star-forming regions."
    },
    "1402/1402.2590_arXiv.txt": {
        "abstract": "Measurements of the cluster abundance as a function of mass and redshift provide an important cosmological test that probe not only the expansion rate but also the growth of perturbations. In this paper we adopt a scalar field scenario which admits both thawing and freezing solutions from an appropriate choice of the model parameters and derived all relevant expressions to calculate the mass function and the cluster number density. We discuss the ability of cluster observations to distinguish between these scalar field behaviors and the standard $\\Lambda$CDM scenario by considering the eROSITA and SPT cluster surveys. ",
        "introduction": "Nowadays, one of the  most important tasks in Cosmology is to unveil the nature of the mechanism responsible for the current cosmic acceleration. From the observational viewpoint, the cosmological constant $\\Lambda$, i.e., the sum over the vacuum energy density of all fields in Nature, remains as the favorite candidate for explaining this phenomenon, but because of its theoretical issues (e.g. fine tuning and cosmic coincidence problems~\\citep{Weinberg}), many alternative mechanisms have been proposed (see \\citep{rev} for recent reviews).  Among these, a very light, minimally coupled scalar field $\\phi$ whose evolution is driven by the slope of the potential $V(\\phi)$ and damped by the cosmic expansion $H$ according to the Klein-Gordon equation $\\ddot{\\phi} + 3H\\dot{\\phi} = - dV/d\\phi $, is the most extensively studied dark energy case~\\citep{quint,Quintessence}. In terms of its potential or still of its equation-of-state (EoS) parameter, $w_\\phi(\\phi, \\dot{\\phi}, V)$, these dark energy fields can be broadly classified into two categories: thawing models whose EoS increases from $w_{\\phi} \\sim -1$, as the field rolls down toward the minimum of its potential with $\\dot w_{\\phi} > 0$ and cooling scenarios in which an initial $w_{\\phi} > -1$ EoS decreases to more negative values with $\\dot w_{\\phi} < 0$. A special case of the latter is the so-called freezing models, in  which the potential has a minimum at $\\phi = \\infty$~\\citep{cl}. Examples of thawing-type scenarios are a pseudo Nambu-Goldstone boson $V(\\phi) = M^4\\cos^2(\\beta \\phi)$~\\cite{ioav} and exponential potentials for dilaton fields $V(\\phi) = M^4\\exp(-\\alpha\\phi)$~\\cite{exp} (see also \\cite{ioav}) whereas the freezing behavior is found in models of the type $V(\\phi) = M^{4 + n}\\phi^{-n}$ and $V(\\phi) = M^{4 + n}\\phi^{-n}\\exp(\\beta \\phi^2)$ for $n >0$~\\cite{freezing}.   Distinguishing among these (and other) physical features of the dark energy potential from cosmological data constitutes an important approach to the cosmic acceleration problem, not only because it may indicate possible routes of solutions but also because it will reduce considerably the range of possibilities. In this regard, measurements of the cluster abundance as a function of mass and redshift are coming up with the potential to improve current constraints on cosmological parameters \\cite{cluster,cluster2}, including the dark energy EoS, the rms mass fluctuations, the matter density parameter, the total neutrino masses, etc. These data have the  advantage of probing both the expansion rate and the growth of perturbations, thereby being complementary to other cosmological probes such as the Comic Microwave background (CMB) anisotropies, observations of type Ia supernovae and measurements of baryon acoustic oscillations (BAO).  In what follows, we investigate to what extent galaxy cluster number counts can be used to distinguish the thawing and freezing behaviors discussed above (see also Sec. II) as well as between these dynamical scenarios and the standard $\\Lambda$CDM model. In our analysis, we consider the scalar field model discussed in Ref.~\\cite{Carvalho2006} which admits both solutions from an appropriate choice of the model parameters. Keeping the range of cosmological parameters within the allowed region by the current observations and considering two specific cluster surveys, namely eROSITA and SPT, we calculate the number density of cluster and the redshift distribution of cluster number counts for this scenario and discuss the ability of these observations to distinguish between the thawing and freezing behaviors. ",
        "conclusions": "Thawing and freezing scalar fields are realistic candidates for dark energy. In terms of the field dynamics, while the former has been frozen until the onset of cosmic acceleration, when it starts to roll down towards the minimum of its potential $V_{min}(\\phi)$, the latter was already rolling to $V_{min}(\\phi)$ at that time and dominates the Universe as it slows down. From the theoretical viewpoint, both behaviors have important implications for the cosmic evolution, with thawing potentials driving the Universe to a transient accelerating phase while freezing models to an eternal quasi-de Sitter expansion.  In this paper we have adopted a specific scalar field model (Eq. 4), which admits both thawing and freezing solutions with an appropriate choice of the parameter $\\alpha$, and discussed the possibility of current and planned survey distinguishing these two behaviors through observations of cluster number counts. Adopting the spherical collapse model, we have studied the structure formation in these scenarios and derived all relevant expressions to calculate the mass function and the cluster number density using the Sheth-Torman formalism. We have found that the cluster counts predictions of the $\\Lambda$CDM and thawing models are very similar, which make this latter class more difficult to be distinguished than freezing ones.  By considering eROSITA and SPT surveys, we have also calculated the theoretically expected cluster redshift distribution for cluster halos size (Fig. 4). From our analysis, we have found that, while SPT would not be able to distinguish between thawing and freezing models, the more precise eROSITA observations have the potential not only to discriminate these two scalar field behaviors but also to differentiate them from the standard $\\Lambda$CDM model."
    },
    "1402/1402.5465_arXiv.txt": {
        "abstract": "We report on a survey for H$\\alpha$ bow shock emission around nearby $\\gamma$-detected energetic pulsars. This survey adds three Balmer-dominated neutron star bow shocks to the six previously confirmed examples. In addition to the shock around {\\it Fermi} pulsar PSR J1741$-$2054, we now report H$\\alpha$ structures around two additional $\\gamma$-ray pulsars, PSR J2030+4415 and PSR J1509$-$5850. These are the first known examples of H$\\alpha$ nebulae with pre-ionization halos. With new measurements, we show that a simple analytic model can account for the angular size and flux of the bow shocks' apices. The latter, in particular, provides a new pulsar probe and indicates large moments of inertia and smaller distances than previously assumed, in several cases. In particular we show that the re-measured PSR J0437$-$4715 shock flux implies $I = (1.7\\pm 0.2) \\times 10^{45}/(f_{HI} {\\rm sin}i) {\\rm g\\,cm^2}$. We also derive a distance $d\\approx 0.72$\\,kpc for the $\\gamma$-ray only pulsar PSR J2030+4415 and revised distances for PSRs J1959+2048 (1.4\\,kpc) and J2555+6535 ($\\sim 1$\\,kpc), smaller than the conventional DM-estimated values. Finally we report upper limits for 94 additional LAT pulsars. An estimate of the survey sensitivity indicates that for a warm neutral medium filling factor $\\phi_{WNM}\\sim 0.3$ there should be a total of $\\sim$ nine H$\\alpha$ bow shocks in our LAT-targeted survey; given that seven such objects are now known, a much larger $\\phi_{WNM}$ seems problematic. ",
        "introduction": "Pulsar wind nebulae (PWNe) have been observed {\\it via} X-ray synchrotron radiation for many sources over many years \\citep[][and references therein]{kp10}. These data probe the bulk energetics of the pulsar wind and, with high quality images from the {\\it Chandra} X-ray observatory, reveal termination shock structure in the form of tori and jets. When the pulsar is subsonic, as in their parent supernova remnant, these X-ray structures can show high symmetry about the spin axis. However, when the pulsar escapes to the cooler external interstellar medium (ISM), the pulsar motion is generally supersonic and the termination shocks are ram pressure confined due to the pulsar motion. This asymmetry can distort the relativistic PWN termination shock. Thus it was particularly interesting when an H$\\alpha$ bow shock was discovered around the black widow millisecond pulsar PSR B1957+20 \\citep{kh88}, both because this structure showed high symmetry about the pulsar velocity axis and because the Balmer line emission allows kinetic study of the shock structure. These bow shocks are non-radiative, with the pulsar velocity ensuring that the shocked ISM flows well down-stream before it cools. Non-radiative shocks show extreme Balmer domination and the velocity structure of these lines turns out to be a sensitive probe of the shock physics, including information on the up-stream ISM, the post-shock flow and even supra-thermal particle acceleration \\citep{waget09}. Pulsar bow shocks have an interesting range of speeds (low compared to most supernova remnants) and shock obliquities, so observation of these nebulae provide an excellent opportunity to study non-radiative shock physics.  Given this interest, a number of searches have been mounted for similar pulsar nebulae \\citep[e.g.][]{crl93,ganet02}, with some limited success. Before the present search there were six confirmed H$\\alpha$ pulsar bow shocks reported in the literature. A shock around PSR J1741$-$2054 was discovered in the initial phase of the present search \\citep{ret10}.  A few other candidates have been presented in the literature, but clearly these objects are quite rare.  We report here on a new sensitive survey of a large number of nearby energetic pulsars. We have tested this sensitivity with images of known bow shocks. Despite the modest exposures, these are in general comparable with or superior to earlier images in the literature, and show additional interesting H$\\alpha$ structure. Our survey has added three new bow shock detections (two reported here for the first time).  We use the new measurements, and a re-assessment of the other known bow shocks, to illustrate the dependence on pulsar properties. In addition we derive quantitative upper limits on the H$\\alpha$ fluxes from the non-detections. This is now a large enough sample that it can be used to evaluate, heuristically, the detection probability's dependence on pulsar parameters. Also, for a basic assumed model, the limits provide an interesting spot sample of the neutral ISM within a few kpc of the Sun. ",
        "conclusions": "In a quest to further understand the rare pulsar H$\\alpha$ bow shocks we have observed a large sample of energetic, {\\it Fermi} LAT-detected pulsars; sensitive re-measurement of known examples provided context for survey imaging of nearly 100 new targets. Three new H$\\alpha$ shocks were discovered, including the detection of two examples of previously unobserved ionization precursors (PSRs J1509$-$5850 and J2030+4415).  To interpret these data we have developed a simple analytic model that describes well the apex zone of all known pulsar bow shocks, including conditions likely to produce precursor halos. This analysis provides a good understanding of the apex zone's angular scale and flux, treating both the high powered pulsars selected by the LAT and the low ${\\dot E}$ nearby pulsars such as PSR J1856$-$3754 and J2225+6535 which apparently feature incomplete ionization in the post-shock flow. This exercise provides a quantitative connection between the bow shock flux and pulsar parameters, applicable over a range of, e.g. $10^6\\times$ in spindown luminosity. We argue (equation 10) that the flux is independent of ISM density and highly sensitive to the pulsar distance. When the distance is poorly known, bow shock observations provide independent measurements of this crucial parameter. We derive a 0.72\\,kpc distance for the $\\gamma$-ray only pulsar J2030+4415. We show that the dispersion measure distances for the millisecond pulsar J1959+2048 and the young, fast PSR J2225+6535 are too large, deriving $d\\approx 1.4$\\,kpc and $d=0.8-1$\\,kpc for these two objects. These distance estimates do depend weakly on velocity inclination $i$, the upstream neutral fraction $f_{HI}$ and, possibly, pulsar wind anisotropy, but these factors can be constrained by measurements of the bowshocks' apex shape and emission line profile. Foreground extinction also affects the estimate, but this can be determined by measuring the nebula Balmer decrement or other absorption studies. Thus pulsar bow shocks, while rare, can provide important new insight into pulsar distances and energetics.  One aspect of this study deserves special attention. The bow show apex covers a large solid angle as viewed from the pulsar and thus provides comprehensive measurement of its mechanical spin-down energy deposition. When a pulsar has a well-measured distance and proper motion, then the bow shock flux puts a lower limit on ${\\dot E}$, which implies a lower limit on the neutron star moment of inertia. PSR J0437$-$4715 is the premier example. Our apex flux of $(6.7\\pm0.7)\\times 10^{-3} {\\rm H\\alpha\\,cm^{-2}\\,s^{-1}}$ translates to a lower limit of $I_{45} > 1.5$ if the pulsar motion is in the plane of the sky and upstream medium is 100\\% neutral.  Since ${\\rm sin}\\,i$ and $f_{HI}$ are $\\le 1$, accounting for these factors will only increase this lower bound. There is in addition a small increase since the photon ($\\gamma$-ray) energy losses do not contribute to the mechanical losses in the wind, but for PSR J0437$-$4715 this accounts for $\\sim 1\\%$ of ${\\dot E}$ (lower luminosity MSP can be substantially more $\\gamma$-efficient). The extinction is small and adds little uncertainty to the luminosity of this well-observed bow shock. We have assumed a spherical pulsar wind, but this can be tested by detailed modeling and measurement of the apex H$\\alpha$ velocity structure. There is one significant assumption in our computation, that the $e^-$ and ions equilibrate (Equation 9); if we use the often-assumed $\\epsilon_{\\rm H\\alpha}=0.2$, our lower bound on the moment of inertia would increase to a (likely unphysical) $I_{45} > 3.9$. If the electrons are completely out of equilibrium the implied moment of inertia is even larger.  Since PSR J0437$-$4715 also has a mass measurement, it should help fulfill the promise of \\citet{ls05} that moment of inertia measurements can pin down the neutron star equation of state. Already our estimated $I_{45}=1.7\\pm 0.2$ prefers relatively stiff EOS. Other pulsar bow shocks, with their near-bolometric measure of the spin-down power can contribute to this quest. It would be very important to obtain a precise independent parallax of PSR J1959+2058, since the high pulsar mass $M_{NS} =2.4M_\\odot$ implies a very high moment of inertia $I_{45} = 2.97$, according to Equation (1). If $I_{45}$ is lower, one would require an even smaller distance than the $d=1.4$\\,kpc determined in this paper. Of course, one or more young neutron stars with low masses and moments of inertia could also be valuable anchors to the equation of state study. Thus precision parallax measurements of the neutron stars showing H$\\alpha$ bow shocks are strongly motivated ingredients in this fundamental physics quest.  Our new images also reveal more details about the post-apex, large scale H$\\alpha$ emission. It is now clear that the nebula limb very generally describes a series of undulations or cavities $10-100\\theta_a$ downstream. While in some cases, these appear nearly closed, other shocks e.g. PSR J0742$-$2822 display oscillation about a more or less conical structure. Thus the guitar nebula is only the most spectacular example of down-stream bubbles. The details doubtless depend on accidents of the ISM, but it seems likely that this behavior is a symptom of instabilities in the post-shock PWN flow \\citep{vKI08}. Surprisingly, we find that with the exception of J1856$-$3754, the total flux of the observed bow shock nebula are very similar, despite their being detected over a wide range of ${\\dot E}$, $v$ and $d$. This calls into question previous efforts to understand the total nebula flux, and suggests that, unlike the apex flux, this does not simply track the pulsar properties.  Finally, our modeling has allowed us to compare our survey flux limits with the expected sizes and fluxes of potential bow shocks around other pulsars.  This study implies that the WNM hosting such bow shocks should occupy no more than $\\sim 30$\\% of the nearby Galactic plane. It appears that additional bow shocks await detection, but that deeper exposures are needed. Our analysis points to a set of high ${\\dot E}$, low extinction objects as the favored targets for future searches, and suggests that intensive study of these nebulae can provide new insights into pulsar physics.  \\medskip  We thank Matt Stadnik for assistance with the image reductions and the referee for a careful reading. This work was supported in part by NASA grants NIX08AW30G and NIX10AD11G and NNX12AO68G, along with Chandra Grant GO0-11097X."
    },
    "1402/1402.7024.txt": {
        "abstract": "The VST Photometric H$\\alpha$ Survey of the Southern Galactic Plane and Bulge (VPHAS$+$) is surveying the southern Milky Way in $u, g, r, i$ and H$\\alpha$ at $\\sim$1~arcsec angular resolution. Its footprint spans the Galactic latitude range $-5^{\\rm o} < b < +5^{\\rm o}$ at all longitudes south of the celestial equator. Extensions around the Galactic Centre to Galactic latitudes $\\pm 10^{\\circ}$ bring in much of the Galactic Bulge.   This ESO public survey, begun on 28th December 2011, reaches down to $\\sim$20th magnitude (10$\\sigma$) and will provide single-epoch digital optical photometry for $\\sim$300 million stars.  The observing strategy and data pipelining is described, and an appraisal of the segmented narrowband $H\\alpha$ filter in use is presented.  Using model atmospheres and library spectra, we compute main-sequence $(u - g)$, $(g - r)$, $(r - i)$ and $(r - H\\alpha)$ stellar colours in the Vega system.  We report on a preliminary validation of the photometry using test data obtained from two pointings overlapping the Sloan Digital Sky Survey.  An example of the $(u-g,g-r)$ and $(r-H\\alpha,r-i)$ diagrams for a full VPHAS+ survey field is given. Attention is drawn to the opportunities for studies of compact nebulae and nebular morphologies that arise from the image quality being achieved. The value of the $u$ band as the means to identify planetary-nebula central stars is demonstrated by the discovery of the central star of NGC 2899 in survey data. Thanks to its excellent imaging performance, the VST/OmegaCam combination used by this survey is a perfect vehicle for automated searches for reddened early-type stars, and will allow the discovery and analysis of compact binaries, white dwarfs and transient sources. ",
        "introduction": "\\label{sec:intro}  The $H\\alpha$ emission line is well-known as a tracer of diffuse ionized nebulae and as a marker of pre- or post-main sequence status among spatially-unresolved stellar sources.  Since these objects -- both nebulae and stars -- represent relatively short-lived phases of evolution, they amount to a minority population in a mature galaxy like our own. Their relative scarcity has in the past stood in the way of developing and testing models for these crucial evolutionary stages.  In the southern hemisphere, the search for planetary nebulae (PNe) has been served well by H$\\alpha$ imaging surveys carried out by the UK Schmidt Telescope (Parker et al 2005, 2006 and other more recent works).  Nevertheless, VPHAS+ will have a decisive impact on studies of complex or smaller nebulae of all types, ranging from optically-detectable ultra-compact and compact HII regions, to nebulae around YSOs (including associated jets and HH objects), through PNe, to extended emission from D-type symbiotic stars and supernova remnants.  The superb spatial resolution, dynamic range, and likely photometric accuracy of the VPHAS+ images warrant a step forward in our knowledge of the population and detailed characteristics of these object classes.  For southern point sources with emission the situation is very different: there has been little updating of the available catalogues since the work of Stephenson \\& Sanduleak (1971) that was limited to a depth of 12th magnitude.  The major groups of emission line stars that remain as challenges to our understanding include all types of massive star (O stars, supergiants, luminous blue variables, Wolf-Rayet stars, various types of Be star), post-AGB stars, pre-main-sequence stars at all masses, active stars and compact interacting binaries.  Within the disc of the Milky Way, the available samples of these objects are typically modest and heterogeneous. Fixing this deficit via a uniform search of the Galactic Plane for these rare object classes motivated the photometric H$\\alpha$ survey of the southern Galactic Plane, first proposed for the VLT Survey Telescope (VST) in 2004.  This paper describes the realisation of this ESO public survey, now known as the VST Photometric H$\\alpha$ Survey of the Southern Galactic Plane and Bulge (VPHAS+).  When first proposed, VPHAS (without the plus sign) was envisaged as the counterpart to the INT/WFC Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS, Drew et al 2005), that had begun in August 2003.  IPHAS is a digital imaging survey made up of Sloan $r$, $i$ and narrowband  $H\\alpha$ exposures, reaching to $\\sim$20th magnitude, that takes in all Galactic longitudes north of the celestial equator in the latitude range $-5^{\\circ} < b < +5^{\\circ}$. This is all but complete with the new release of a catalogue of $\\sim$200 million unique objects drawn from 93 percent of the survey's footprint (Barentsen et al, 2014). During the initial VST public survey review process, it was agreed that VPHAS should broaden in scope to also incorporate the Sloan $u$ and $g$ bands (proposed for a separate survey), that are particularly useful in picking out OB stars, white dwarfs and other blue-excess objects.  With this upgrade to 5 bands, the renamed VPHAS$+$ became an all-purpose digital optical survey of the southern Galactic Plane, capable of delivering data at a spatial resolution of $\\sim$1~arcsec or better.  As well as fulfilling the role of southern counterpart of IPHAS, VPHAS$+$ is also the counterpart to UVEX, the UV-Excess Survey of the Northern Galactic Plane (Groot et al, 2009) that, at the time of writing, continues on the Isaac Newton Telescope in La Palma.  The final augmentation of the VPHAS$+$ survey footprint came in 2010 on expanding its footprint to match that of the similarly high spatial-resolution near-infrared survey, VISTA Variables in the Via Lactea (VVV, Minniti et al 2011).  The survey footprint now includes the Galactic Bulge to a latitude of $|b| < 10^{\\circ}$, across the longitude range $-10^{\\circ} < \\ell < +10^{\\circ}$.  The VVV $z$, $Y$, $J$, $H$ and $K_s$ survey of much of the Bulge and inner Galactic disc is already complete.  VPHAS$+$ is poised to become the homogeneous digital optical imaging survey of the Galactic Plane and Bulge, at $\\sim$1~arcsec angular resolution, that will provide a uniform database of stellar spectral energy distributions, from which a range of colour-magnitude and colour-colour diagrams of well-established utility can be derived. Once calibrated to the expected precision of 2 to 3 percent, like its northern counterparts, VPHAS$+$ will be quantitatively far superior to the photographic surveys of the last century, and will offer significant added value in the form of calibrated narrowband H$\\alpha$ data.  Some of the science enabled is illustrated by the studies that the northern surveys, IPHAS and UVEX, have already stimulated.  These have included a number of works reporting the discovery of emission line stars, ranging from young objects (e.g. Valdivieso et al, 2009, Vink et al 2008, Barentsen et al 2011, Raddi et al 2013) to evolved object classes such as symbiotic and cataclysmic binaries and compact planetary nebulae (e.g. Corradi et al 2010, Witham et al 2006, Viironen et al 2009 and Wesson et al 2008). The diagnostic power to be expected from the blue $u$ and $g$ bands has been appraised in a series of papers presenting UVEX and early follow-up spectroscopy (Groot et al 2009, Verbeek et al 2012a and 2012b).  \\begin{figure*} \\includegraphics[angle=0,width=\\linewidth,trim=0 0 0 500]{vphas_fig1.pdf} %{\\special{psfile=vphas_fig1.eps %angle =0 hoffset=-230 voffset=-5 hscale=88 vscale=88}} %\\end{picture} \\caption{The VPHAS$+$ survey footprint plotted in Galactic coordinates. All 2269 fields are shown in outline.  Different colours, as specified in the key, are used to identify the observations obtained for each field by 1st January 2014 -- essentially 2 years after the start of data-taking.} \\label{fig:footprint} \\end{figure*}  Any comprehensive survey of the Galactic Plane clearly targets the main mass component of our own Galaxy, made up of stars, gas and dust.  When a spatial resolution of 1 arcsec is combined with wide area coverage spanning 100s of square degrees, as in IPHAS, UVEX and VPHAS$+$, it becomes possible to exploit the stellar photometry adaptively to solve for the distribution of both the stars and dust making up the optically-accessible Galactic disk - by means that are similar to those already attempted, based on the near-infrared 2MASS survey (Drimmel \\& Spergel 2001, Marshall et al 2006). Methods to achieve 3-dimensional mapping of this kind, now incorporating the direct sensitivity of the H$\\alpha$ narrowband to stellar intrinsic colour, are starting to take shape (Sale et al 2009; Sale 2012). This development coincides with the approach of the operations phase of Europe's next major astrometric mission, Gaia.  Indeed, the rich dynamical picture that Gaia will build of the Milky Way over the next decade will be very effectively complemented by stellar energy distributions measured for millions of stars from the current generation of ground-based optical, and near-infrared, wide-field surveys.  VPHAS$+$, with its haul of photometry in 5 optical bands on $\\sim$300 million objects, is set to take its place as one of them.  %This last group %most likely harbours SN~Ia progenitors within it, in guises that are still %subject to considerable debate (Hillebrandt \\& Niemeyer 2000; Uenishi, Nomoto %\\& Hachisu 2003).  This paper presents the main features of VPHAS$+$, including a description of its execution, the data processing and the nature of the photometric colour information it provides.  We begin in the next section with a presentation of the observing strategy and the data reduction techniques in use.  We then turn to a description and evaluation of the narrowband H$\\alpha$ filter procured for this survey, in section~\\ref{sec:ha}.  Following this, in section~\\ref{sec:simul}, we present tailored synthetic photometry of main sequence and giant stars that provides insights into the photometric diagrams that may be generated from survey data.  An exercise in photometric validation is described in Section~\\ref{sec:phot-val} in which Sloan Digital Sky Survey (SDSS) data are compared with VST observations.  The scene is then set for an example of VPHAS+ photometry extracted across the entire square-degree footprint of a single survey field (Section~\\ref{sec:example}). In Sections~\\ref{sec:imaging} and \\ref{sec:var}, we outline the applications of VPHAS$+$ to spatially-resolved nebular astrophysics and in the time domain.  The paper ends in Section~\\ref{sec:last} with a summary, examples of early survey exploitation, and a forward look to the first major data release. ",
        "conclusions": "\\label{sec:last}  This paper has introduced and defined VPHAS$+$, the VST Photometric $H\\alpha$ Survey of the Southern Galactic Plane and Bulge. The data taking, the rationale behind it, the data processing and data quality have all been described.  The properties and limitations of the survey's narrowband H$\\alpha$ filter, NB-659, have been laid out and simulated in order to anticipate its performance. In addition we have provided tables of the expected photometric colours of normal solar-metallicity stars to aid the interpretation of the survey's characteristic photometric diagrams -- most are to be found in the Appendix where the effect of changing the adopted reddening law is illustrated.  The VPHAS+ H$\\alpha$ filter transmission is redder, wider and $\\sim20$ \\% higher-throughput than its IPHAS counterpart -- a difference that feeds through to noticeably different $(r - H\\alpha)$ colours for M stars.  We have validated the photometry that is delivered by VST/OmegaCAM and subsequently pipelined at CASU, using test data taken of a field for which SDSS photometry is available. We find the agreement is satisfactory, with the $g$, $r$ and $i$ band calibrations differing by between 0.01 to 0.05 magnitudes.  However, for the time being, the pipeline calibration should be regarded as provisional -- it will undoubtedly improve.  Examples of the excellent imaging performance of the VST/OmegaCAM combination relative to previous surveys have been provided, and we draw attention to the valuable archival role this first digital survey can fulfill in supporting discoveries of transient sources.  Exploitation of the survey is now beginning.  The detection of a compact ionized nebula around W26, the extreme M supergiant in Westerlund 1 has already been published (Wright et al 2014).  Applications have been made for follow-up spectroscopy that will test the quality of selection of specialised object types that VPHAS+ photometry makes possible. Progress is also being made via direct analyses of the photometry. For example, Mohr-Smith et al (in prep.) are conducting a search for OB stars in the vicinity of the massive cluster Westerlund 2, and they are finding a close match between the properties of known cluster O stars as derived from VPHAS+ data and those inferred by Vargas Alvarez et al (2013, see also Drew et al 2013). This and other early appraisals of the data indicate that VPHAS+ will be an excellent vehicle for automated searches for reddened early-type stars.  Kalari et al (in prep.) are employing both narrow-band H$\\alpha$ and the broadbands to measure mass-accretion rates in pre-main-sequence stars: they are finding that H$\\alpha$ mass-accretion rates in T~Tauri stars compare favourably to rates determined from the $u$ band in the case of the Lagoon Nebula, NGC~6530.  As the calibration of the survey data improves, the measurement of accurate integrated H$\\alpha$ fluxes for many faint southern PNe and other extended objects becomes possible, and will extend the work of Frew et al (2013, 2014).  In due course these fluxes can be compared with existing and also new radio continuum fluxes coming on stream (see e.g. Norris et al 2011) in order to determine reliable extinction values for many faint nebulae currently lacking data.  This technique has already been applied to the case of W26 in Westerlund 1 (Wright et al 2014).  When it becomes possible to cross-match VVV and VPHAS$+$ data, it will open up the power of homogeneous photometric mapping of the central parts of the Galactic Plane in up to 10 photometric bands spanning both the optical and the near-infrared. Beyond the VVV sky area, there is a synergy to be exploited in bringing VPHAS+ data together with those of the all-sky 2MASS survey (Skrutskie et al 2006) and with the UKIDSS Galactic Plane Survey (Lucas et al 2008), in those parts of the first and third Galactic quadrants the latter has covered. It is worth noting, however, that 2MASS alone is too shallow to link effectively with VPHAS+ for sightlines where the integrated visual extinction is less than $\\sim 5$ magnitudes.  This does mean that the longitude range $230^{\\circ} < \\ell < 300^{\\circ}$, in particular, is presently lacking sufficiently deep NIR photometry.  In the longer term, many of the sources of interest that VPHAS+ finds will benefit from accurate parallaxes and other data from ESA's Gaia mission -- given the similar sensitivity limits reached.  Conversely in the meantime, VPHAS+ has already begun to assist ambitious wide-field spectroscopy programmes such as the Gaia-ESO Survey (Gilmore et al 2012) through the provision of the wide-field photometry needed for target selection and field setup.  By the end of 2013, 25\\% of all observations making up the survey had been obtained to the required quality, and in May 2013 a first release of single-band catalogues was made to the ESO archive that contained roughly 10\\% of the eventual total (based on data obtained prior to 15 October 2012).  By design, the characteristics of VPHAS+ are similar to those of the IPHAS and UVEX Galactic plane survey pair in the north.  In particular, the double-pass strategy is shared, with the result that the majority of detected objects are picked up and measured twice, with no more than $\\sim$0.1 percent of objects missed altogether.  This feature has informed the way in which the IPHAS DR2 catalogue (Barentsen et al, 2014) has been constructed -- and it is intended that a first band-merged VPHAS+ catalogue, for public release, will be built along analogous lines during the second half of 2014.  This will incorporate data from the first 3 seasons of VST observing, and give a complete photometric account of the Galactic mid-plane.  For ease of use, for every detected source, the catalogue will provide a single recommended set of magnitudes in up to 5 optical bands."
    },
    "1402/1402.1520_arXiv.txt": {
        "abstract": "Recent solution of a pulsar (although likely correct, as agreeing with phenomenology of weak pulsars without a single adjustable) is somewhat unsatisfactory on account of being heavily numerical. Here the main features of a pulsar -- (i) the existence of a force-free zone, (ii) bounded by a non-force-free radiation zone, (iii) where colliding  Poynting fluxes annihilate into curvature radiation -- are reproduced as an exact solution of the same pulsar equations but in a much simpler geometry. ",
        "introduction": "We first give a qualitative description of a pulsar according to \\cite{Gruzinov2013}, \\S\\ref{pic}. This description is based on numerical solutions of Aristotelian Electrodynamics (AE = Electrodynamics of Massless Charges). AE equations are summarized in \\S\\ref{AE}. In \\S\\ref{dev} we give an exact solution of AE equations for ``the Device'' which was specially chosen to be exactly solvable and similar to real pulsars. Namely, the Device features: (i) a force-free zone, (ii) bounded by a non-force-free radiation zone, (iii) where colliding  Poynting fluxes annihilate into curvature radiation. ",
        "conclusions": "\\begin{itemize} \\item Pulsars shine by annihilating colliding Poynting Fluxes into curvature radiation. \\item To solve a pulsar one can use AE. \\item The Device, being exactly solvable in the AE weak-pulsar limit, can prove useful in testing  numerical codes, especially when the weak-pulsar and/or AE approximations are lifted. \\end{itemize}"
    },
    "1402/1402.4133_arXiv.txt": {
        "abstract": "{Protoplanetary discs are poorly ionised due to their low temperatures and high column densities, and are therefore subject to three ``non-ideal'' magnetohydrodynamic effects: Ohmic dissipation, ambipolar diffusion, and the Hall effect. The existence of magnetically driven turbulence in these discs has been a central question since the discovery of the magnetorotational instability. Early models considered Ohmic diffusion only and led to a scenario of layered accretion, in which a magnetically ``dead'' zone in the disc midplane is embedded within magnetically ``active'' surface layers at distances $\\sim$1--10 au from the central protostellar object. Recent work has suggested that a combination of Ohmic dissipation and ambipolar diffusion can render both the midplane and surface layers of the disc inactive and that torques due to magnetically driven outflows are required to explain the observed accretion rates. We reassess this picture by performing three-dimensional numerical simulations that include, for the first time, all three non-ideal MHD effects. We find that the Hall effect can generically ``revive'' dead zones by producing a dominant azimuthal magnetic field and a large-scale Maxwell stress throughout the midplane, provided the angular velocity and magnetic field satisfy $\\bm{\\Omega\\cdot B} > 0$. The attendant large magnetic pressure modifies the vertical density profile and substantially increases the disc scale height beyond its hydrostatic value. Outflows are produced, but are not necessary to explain accretion rates $\\lesssim$$10^{-7}~{\\rm M}_\\odot~\\mathrm{yr}^{-1}$. The flow in the disc midplane is essentially laminar, suggesting that dust sedimentation may be efficient. These results demonstrate that, if the MRI is relevant for driving mass accretion in protoplanetary discs, one must include the Hall effect to obtain even qualitatively correct results.} ",
        "introduction": "The formation and evolution of protoplanetary discs is a long-standing problem in accretion theory. Observations indicate that T Tauri stars accrete mass from their surrounding discs at a rate $\\sim$ $10^{-9}$--$10^{-7}~\\mathrm{M_\\odot~yr}^{-1}$ on timescales $\\sim$$1$--$10~{\\rm Myr}$ \\citep{HC98}. However, the physical mechanism driving this accretion has remained elusive for decades. Such a mechanism not only is necessary for understanding how the central protostellar object accretes mass at the observed rates, but also represents a crucial component in any theory for the formation of planets. In order to effectively accrete, angular momentum must be removed from the infalling gas and transferred to other fluid elements within the disc and/or in the surrounding medium. It can either be \\emph{transported radially}, in which case the transport is usually modelled as a viscous process (so-called $\\alpha$-disc models; \\citealt{SS73}), or \\emph{transported vertically} by large-scale magnetically driven outflows or winds \\citep{BP82}.  The rediscovery of the magnetorotational instability (MRI; \\citealt{V59,C60}) by \\citet{BH91,BH98} in the context of astrophysical discs vastly improved our understanding of enhanced angular-momentum transport in differentially rotating plasmas. This instability produces vigorous turbulence and efficiently transports angular momentum outwards in discs with an effective viscosity $\\alpha \\gtrsim 10^{-3}$ \\citep{HGB95}. The MRI has since become the main contender\\footnote{Hydrodynamical processes are an interesting alternative, despite serious theoretical difficulties; see \\cite{T14} for a review.} to explain accretion in many types of astrophysical systems. However, protoplanetary discs are cold ($T\\lesssim 300~K$ at 1 au), dense ($n\\gtrsim 10^{13}~\\mathrm{cm}^{-3}$), and thus poorly ionised. Their evolution is well-described  by a set of magnetohydrodynamic (MHD) equations that include three non-ideal effects \\citep{NH85,WN99}: Ohmic dissipation, ambipolar diffusion, and the Hall effect. In a magnetised plasma composed of neutral molecules, ions, and electrons, Ohmic dissipation is caused by collisions between electrons and neutrals, ambipolar diffusion by collisions between ions and neutrals, and the Hall effect by the velocity difference (``drift'') between electrons and ions (or, more generally, between positively and negatively charged species). These non-ideal effects tend to decouple the gas from the magnetic field, casting doubt upon the relevance of the MRI to protoplanetary discs \\citep{BB94,W99,KB04,D04}.  The impact of Ohmic dissipation on stratified protoplanetary discs was first studied by \\citet{G96} and \\citet{SM99}, and led to a model of layered accretion at distances $\\sim$1--10 au from the central protostellar object: a quiescent magnetically ``dead'' midplane is surrounded by turbulent magnetically ``active'' surface layers located at a few disc scale heights. This layered-accretion model has been studied extensively with increasingly complex analytical models \\citep{SMUN00,FTB02,IN06a,IN06b} and three-dimensional numerical simulations \\citep{FSH00,FS03,TSD07,IN08}. Depending upon the chemical network and the presence of well-mixed dust grains, these models predict a vertically integrated turbulent transport roughly consistent with the observed accretion rates. However, recent work by \\citet{BS13} and \\citet{SB13} has shown that ambipolar diffusion can render the surface layers magnetically inactive. Magnetically driven outflows, which naturally occur in stratified shearing boxes including a mean vertical field \\citep{SI09,M12,O12,LFO13,FLLO13,BS13a}, are then needed to produce accretion rates comparable to those observed.  Although the influence of the Hall effect on the linear properties of the MRI is well understood \\citep{W99,BT01}, its impact on the non-linear evolution, and in particular on the dynamics of putative dead zones, has remained mostly speculative \\citep{SS02,SW03,SW05,W07,SW08,WS12}. In a recent publication, \\citet[][hereafter KL13]{KL13} discovered a new saturation mechanism for Hall-dominated magnetorotational turbulence, whereby the vertical magnetic field self-organises into coherent, long-lived, axisymmetric (``zonal'') structures and the rate of angular-momentum transport becomes vanishingly small. While interesting in itself, this result was based on a crude representation of a protoplanetary disc -- an incompressible, dust-free, unstratified shearing box without ion-neutral drift. It is natural to ask whether this result applies to stratified models including Ohmic dissipation and ambipolar diffusion.  This is precisely the question we address in this paper. We perform three-dimensional numerical simulations of the MRI in a compressible, vertically stratified shearing box, taking into account Ohmic dissipation, ambipolar diffusion, and the Hall effect. This marks the first time that all three non-ideal magnetohydrodynamic (MHD) effects have been accounted for in a numerical simulation of the MRI. To accomplish this feat, we have implemented an original formulation of ambipolar diffusion and the Hall effect in the finite-volume code \\textsc{Pluto}. A simplified ionisation model is used to determine vertical diffusivity profiles for all three non-ideal MHD effects at distances of 1, 5, and 10 au from the central protostellar object.  The paper is organised as follows. In Section \\ref{sec:sb} we present the equations of non-ideal MHD in the stratified shearing-box approximation. In Sections \\ref{sec:num_method} and \\ref{sec:diagnostics} we introduce our numerical method and define the diagnostics used to quantify the efficiency of angular-momentum transport and the production of outflows. We defer to an Appendix the technical details and tests of our numerical implementations of the Hall effect and ambipolar diffusion. We then present the model used to determine the ionisation rate and the diffusivity tensor (\\S\\ref{sec:ionisation}). The results are presented, first with models that neglect ambipolar diffusion (\\S\\S\\ref{sec:linHall}--\\ref{sec:nlHall}) and then with all three non-ideal MHD effects treated self-consistently (\\S\\ref{sec:themeat}). Since our ionisation model does not account for the chemical and dynamical effects of dust grains, we briefly address their impact on our results in Section \\ref{sec:grains}. Finally, in Section \\ref{sec:conclusions} we provide a summary of our results and comment on their implications for the evolution of protoplanetary discs and the formation of planetesimals. ",
        "conclusions": "the Hall effect remains the dominant non-ideal term in disc midplanes at 5 and 10 au, with diffusivity coefficients comparable to that at 1 au without grains. 0.1$\\mum$-sized grains have a more severe impact on the diffusivities, though Hall remains an important player at 5 au. We have not explored this regime numerically because of the cost of our explicit numerical scheme. We note, however, that there is no reason \\emph{a priori} to neglect the Hall effect at these distances.  While a more sophisticated treatment of the chemistry -- in which dust grains are included and species abundances are time-dependent -- is necessary to make accurate quantitative predictions, the results presented here provide unequivocal evidence that Ohmic dissipation, ambipolar diffusion, and the Hall effect must all be taken into account to obtain a realistic description of angular-momentum--transport in protoplanetary discs."
    },
    "1402/1402.4243_arXiv.txt": {
        "abstract": "Cherenkov telescopes have the capability of detecting high energy tau neutrinos in the energy range of 1--1000\\,PeV by searching for very inclined showers. If a tau lepton, produced by a tau neutrino, escapes from the Earth or a mountain, it will decay and initiate a shower in the air which can be detected by an air shower fluorescence or Cherenkov telescope. In this paper, we present detailed Monte Carlo simulations of corresponding event rates for the VERITAS and two proposed Cherenkov Telescope Array sites: Meteor Crater and Yavapai Ranch, which use representative AGN neutrino flux models and take into account topographic conditions of the detector sites. The calculated neutrino sensitivities depend on the observation time and the shape of the energy spectrum, but in some cases are comparable or even better than corresponding neutrino sensitivities of the IceCube detector. For VERITAS and the considered Cherenkov Telescope Array sites the expected neutrino sensitivities are up to factor 3 higher than for the MAGIC site because of the presence of surrounding mountains. ",
        "introduction": "Neutrinos have long been anticipated to help answering some fundamental questions in astrophysics like the mystery of the source of the cosmic rays (for a general discussion see \\cite{uhecr}). For neutrinos in the TeV range, prime source candidates are Galactic supernova remnants \\cite{remnats}. Neutrinos in the PeV range and above are suspected to be produced by Active Galactic Nuclei (AGN) and Gamma Ray Bursts (GRB) with many AGN models predicting a significant neutrino flux \\cite{Atoyan:2004pb,PhysRevD.80.083008,Mucke2003593}. Recently, the IceCube Collaboration has reported the very first observation of a cosmic diffuse neutrino flux which lies in the 100 TeV to PeV range~\\cite{science_icecube}. Individual sources, however, could not be identified up to now. While many astrophysical sources of origin have been suggested \\cite{icecube_inter}, there is yet not enough information to narrow down the possibilities to any particular source.  Due to the low interaction probability of neutrinos, a large amount of matter is needed in order to detect them. One of the detection techniques is based on the observation of inclined extensive air showers (EAS) induced by taus from tau neutrino interactions deep in the atmosphere. As these showers are initiated close to the surface of the Earth, they are still very young when reaching the detector and hence have a significant electromagnetic component leading to a broad time structure of the detected signal. In contrast, showers from cosmic-ray nuclei are induced in the upper atmosphere and therefore have a strongly reduced electromagnetic component when reaching the detector. However, because of its low density, neutrino interactions are not very likely to happen inside the atmosphere. A solution to this problem is to look for so-called Earth skimming (up-going) tau neutrinos  \\cite{fargion0,fargion1,zas,Letessier:2001,feng,tseng,aramo} which interact within the Earth or a mountain and produce a tau. For neutrino energies of about a EeV, the charged leptons have a range of a few kilometers and hence may emerge from the Earth or mountain, decay shortly above the ground and produce EAS detectable by a surface detector. In some cases, two consecutive EASs might be observable, one coming from a tau neutrino interaction close to the surface and one from the decay of the resulting tau lepton. These two showers, coming from the same direction in a time interval corresponding to the tau decay time, can generate a unique signature in the detector called a Double-Bang event ~\\cite{learne}. The detection of such a Double-Bang event would be very important both from the astrophysical and the particle physics point of view, as it would be an unambiguous sign for an ultra-high energy (UHE) tau-neutrino. Up to now, there has been no clear identification of tau neutrinos at high energies.  The detection of PeV tau neutrinos through optical signals also seems possible. A combination of fluorescence and Cherenkov light detectors in the shadow of steep cliffs could achieve this goal  {\\bf ~\\cite{fargion0,fargion1,doi:10.1142/S0217732304014458}}. Recently, it has also been shown by the All-sky Survey High Resolution Air-shower detector (Ashra) experiment, that such kind of experiments could be sensitive to tau neutrinos from fast transient objects such as nearby GRBs~\\cite{Asaoka:2012em}.We note, that the recent IceCube results do not show any neutrino events at or above the Glashow resonance at 6.3 PeV~\\cite{science_icecube}. This likely means that there is either a cutoff in the astrophysical neutrino flux below $\\sim 6$ PeV or the neutrino spectrum is steeper than the usually assumed $E^{-2}$ spectrum.  In principle also existing Imaging Air Cherenkov Telescopes (IACTs) such as MAGIC~\\cite{magic}, VERITAS~\\cite{veritas} and H.E.S.S.~\\cite{hess} have the capability to detect PeV tau neutrinos by searching for very inclined showers~\\cite{fargion}. In order to do that, the Cherenkov telescopes need to be pointed in the direction of the taus escaping from the Earth crust, i.e.\\ at or a few degree below the horizon. This is because the trajectory of the tau lepton has to be parallel to the pointing direction of the telescope within a few degrees as the Cherenkov light is very much beamed in the forward direction. For example, the MAGIC telescope is placed on top of a mountain on La Palma at an altitude of about 2200\\,m a.s.l. Thus, it can look down to the Sea and monitor a large volume within its field of view (FOV). In \\cite{upgoing_magic}, the effective area for up-going tau neutrino observations with the MAGIC telescope was calculated analytically with the maximum sensitivity in the range from 100\\,TeV to $\\sim$~1\\,EeV. However, the calculated sensitivity for diffuse neutrinos was very low because of the limited FOV, the short observation time and the low expected neutrino flux.  On the other hand, if flaring or disrupting point sources such as GRBs are observed, one can  expect an observable number of events even from a single GRB if close by. In the case of MAGIC, however, the topographic conditions allow only for a small window of about 1 degree width in zenith and azimuth to point the telescope downhill. In case of other IACT sites with different topographic conditions, the acceptance for up-going tau neutrinos will be increased by the presence of mountains. Mountains can work as an additional target and will lead to an enhancement in the flux of emerging tau leptons. A target mountain can also shield against cosmic rays and star light.  For Cherenkov telescope sites, very often nights with high clouds prevent the observation of gamma-ray sources. In such conditions, pointing the telescopes to the horizon could significantly increase the observation time and the acceptance for up-going tau neutrinos. Next-generation Cherenkov telescopes, i.e.\\ the Cherenkov Telescope Array (CTA)~\\cite{cta}, can in addition exploit their much larger FOV (in extended observation mode), a higher effective area and a lower energy threshold.  In this work, we present an update of the work in~\\cite{ouricrc}, where a detailed Monte Carlo simulation of event rates induced by Earth skimming tau neutrinos was performed for an ideal Cherenkov detector. Neutrino and lepton propagation was simulated taking into account the local topographic conditions at the MAGIC site and at four  possible locations of Cherenkov instruments: two in Argentina (San Antonio, El Leoncito), one in Namibia (Kuibis) and one on the Canary Islands (Tenerife). In this work, similar simulations have been performed for the location of the VERITAS telescopes and for two sites located close to VERITAS: Meteor Crater and Yavapai Ranch. These two sites were recently also considered as possible locations for CTA. Results are shown for a few representative neutrino fluxes expected from giant AGN flares. We would like to stress that in this work we are exploring the effect of different topographic conditions rather than providing a comprehensive survey of potential sites. ",
        "conclusions": "In this paper, detailed Monte Carlo simulations have been employed to evaluate the sensitivity of air Cherenkov telescopes to high-energy  cosmic tau neutrinos in the PeV to EeV energy range coming from the horizon. The simulations use recent predictions for neutrino fluxes in AGN flares to calculate expected event rates at the VERITAS  site and two  future locations of Cherenkov instruments: Meteor Crater and Yavapai Ranch, taking into account the local topographic conditions. The calculated neutrino rates are comparable to what has been estimated for the IceCube neutrino telescope assuming realistic observation times for Cherenkov telescopes of a few hours. For models with significant neutrino fluxes at energies above $\\sim  5\\times 10^{17}$\\,eV, the sensitivity of Cherenkov telescopes actually surpasses that of IceCube. For the sites considered, the expected event rates are up to factor three higher compared to what is expected for the La Palma site (MAGIC)."
    },
    "1402/1402.4075_arXiv.txt": {
        "abstract": "{} {We analyze the applicability of far-infrared fine-structure lines [C{\\sc{ii}}] 158\\,$\\mu$m, [O{\\sc{i}}] 63\\,$\\mu$m and [O{\\sc{iii}}] 88\\,$\\mu$m to reliably trace the star formation rate (SFR) in a sample of low-metallicity dwarf galaxies from the \\textit{Herschel} Dwarf Galaxy Survey and compare with a broad sample of galaxies of various types and metallicities in the literature.} {We study the trends and scatter in the relation between the SFR (as traced by GALEX $FUV$ and MIPS\\,24\\,$\\mu$m) and far-infrared line emission, on spatially resolved and global galaxy scales, in dwarf galaxies. We assemble far-infrared line measurements from the literature and infer whether the far-infrared lines can probe the SFR (as traced by the total-infrared luminosity) in a variety of galaxy populations.} {In metal-poor dwarfs, the [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ lines show the strongest correlation with the SFR with an uncertainty on the SFR estimates better than a factor of 2, while the link between [C{\\sc{ii}}] emission and the SFR is more dispersed (uncertainty factor of 2.6). The increased scatter in the SFR-$L_{\\text{[CII]}}$ relation towards low metal abundances, warm dust temperatures and large filling factors of diffuse, highly ionized gas suggests that other cooling lines start to dominate depending on the density and ionization state of the gas. \\\\ For the literature sample, we evaluate the correlations for a number of different galaxy populations. The [C{\\sc{ii}}] and [O{\\sc{i}}]$_{63}$ lines are considered to be reliable SFR tracers in starburst galaxies, recovering the star formation activity within an uncertainty of factor 2. For composite and AGN sources, all three FIR lines can recover the SFR with an uncertainty factor of 2.3. The SFR calibrations for ULIRGs are similar to starbursts/AGNs in terms of scatter but offset from the starburst/AGN SFR relations due to line deficits relative to their total-infrared luminosity. While the number of detections of the FIR fine-structure lines is still very limited at high-redshift for [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$, we provide a SFR calibration for [C{\\sc{ii}}]. } {} ",
        "introduction": "Star formation encompasses the birth of new stars through the fragmentation and contraction of cold, dense molecular gas, hereby recycling the ISM content of galaxies \\citep{2007ARA&A..45..565M}. Knowing the instantaneous level of star formation (i.e. the star formation rate) not only sheds light on the conditions in the ISM, but also on the evolution of galaxies and their formation processes. If we want to understand the physical processes that control galaxy evolution, being able to probe the star formation activity is of great importance. Tracing back star formation in some of the first objects in the early Universe until the present-day could even allow us to probe the star formation activity throughout cosmic times.  Since the main coolants in metal-rich galaxies such as the Milky Way are mostly metal-based ([C{\\sc{ii}}], [O{\\sc{i}}], [O{\\sc{iii}}], CO, dust), the metal abundance is considered a fundamental parameter in the regulation of star formation through its influence on the initial cooling of diffuse gas (e.g. \\citealt{2014MNRAS.437....9G}) and the survival of clouds through shielding. Star formation in the early Universe is considered to differ significantly from the present day's gas consumption in galaxies. Due to the extremely low metal abundances in the early Universe, the gas coolants and initiation processes of star formation were likely to be different from star formation conditions in the Local Universe (e.g. Ly $\\alpha$ cooling becomes more important). Nearby low-metallicity galaxies might be important laboratories to investigate the connection between the chemical enrichment and star formation processes. Although the present-day metal-poor dwarfs will have experienced some evolution throughout cosmic time (e.g. \\citealt{1991ApJ...379..621H,2007IAUS..235...65T}), their slow chemical evolution makes them important testbeds to understand metallicity effects potentially applicable to galaxies in the early Universe.  Star formation rates (SFRs) on global galaxy scales are typically estimated from scaling relations between diagnostic tracers of the star formation activity, calibrated against the most up to date stellar population synthesis models and characterization of the initial mass function (IMF). Up to the present day, the most widely used SFR diagnostics have been based on continuum bands and optical/near-IR emission lines (see \\citealt{1998ARA&A..36..189K} and \\citealt{2012ARA&A..50..531K} for a detailed overview). The brightest cooling lines in the atomic and molecular medium are emitted from mid-infrared to radio wavelengths, which have been probed extensively with the \\textit{Herschel} Space Observatory \\citep{2010A&A...518L...1P}. Follow-up is guaranteed with SOFIA in the local Universe or, at high spatial resolution, with ground-based interferometers such as ALMA and the future NOEMA in the high-redshift Universe. These facilities open up a whole new spectral window, which favors the use of far-infrared (FIR) and submillimeter (submm) continuum bands (e.g. PACS\\,70\\,$\\mu$m, \\citealt{2013ApJ...768..180L}) and emission lines as SFR diagnostics across a large variety of galaxy populations.  Here, we investigate the utility of the three brightest fine-structure cooling lines, [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ (e.g. \\citealt{2001ApJ...553..121H,2008ApJS..178..280B}; Cormier et al. in prep.), as tracers of the star formation activity in a sample of low-metallicity dwarf galaxies from the \\textit{Herschel} Dwarf Galaxy Survey \\citep{2013PASP..125..600M}. We, furthermore, asses the influence of metallicity, which is an important parameter controlling star formation in galaxies, constraining the reservoir of dust grains for the formation of molecules and regulating the attenuation of the FUV photons necessary for shielding molecules.  The [C{\\sc{ii}}] 157.74 $\\mu$m line has been put forward as a potential powerful tracer of the star formation activity in the nearby as well as the more distant Universe \\citep{1991ApJ...373..423S,2002A&A...385..454B,2010ApJ...724..957S,2011MNRAS.416.2712D,2012ApJ...755..171S} and we, now, aim to expand this analysis to nearby, low-metallicity dwarf galaxies. [C{\\sc{ii}}] is considered to be the dominant coolant for neutral atomic gas in the interstellar medium \\citep{1985ApJ...291..722T,1985ApJ...291..747T,1995ApJ...443..152W} and, therefore, among the brightest emission lines originating from star-forming galaxies (e.g. \\citealt{1991ApJ...373..423S,1997ApJ...491L..27M,2008ApJS..178..280B}). In particular, low-metallicity galaxies show exceptionally strong [C{\\sc{ii}}] line emission (e.g. \\citealt{1995ApJ...454..293P}, \\citealt{1997ApJ...483..200M}, \\citealt{2000NewAR..44..249M}, \\citealt{2001ApJ...553..121H}, \\citealt{2010A&A...518L..57C}, \\citealt{2011A&A...531A..19I}). Carbon has an ionization potential of 11.3 eV (compared to 13.6 eV for hydrogen), implying that line emission can originate from neutral and ionized gas components (see Table \\ref{critical}). A changing contribution of different gas phases on global scales can prevent a correlation between the [C{\\sc{ii}}] line emission and level of star formation. The excitation of C$^{+}$ atoms might, furthermore, saturate at high temperatures, where the line becomes insensitive to the intensity of the radiation field at temperatures well above the excitation potential \\citep{1999ApJ...527..795K}. The [C{\\sc{ii}}] emission can also saturate in neutral gas media with hydrogen densities $n_{\\text{H}}$ $\\gtrsim$ 10$^{3}$ cm$^{-3}$, where the recombination of C$^{+}$ into neutral carbon and, eventually, CO molecules is favored \\citep{1999ApJ...527..795K}. Self-absorption can also affect the [C{\\sc{ii}}] line excitation in large column densities of gas ($N_{\\text{H}}$ $\\sim$ 4 $\\times$ 10$^{22}$ cm$^{-2}$, \\citealt{1997ApJ...491L..27M}). Although the [C{\\sc{ii}}] line is usually not affected by extinction, optical depth effects might become important in extreme starbursts \\citep{1998ApJ...504L..11L,2000isat.conf..337H} and edge-on galaxies \\citep{1994ApJ...436..720H}. On top of this, deficits in the [C{\\sc{ii}}]/FIR ratio towards warm dust temperatures \\citep{1985ApJ...291..755C,1991ApJ...373..423S,1997ApJ...491L..27M,2001ApJ...561..766M,2003ApJ...594..758L,2005SSRv..119..355V,2008ApJS..178..280B,2011ApJ...728L...7G,2012ApJ...747...81C,2013ApJ...774...68D,2013ApJ...776...38F} suggest that the [C{\\sc{ii}}] conditions might be different in galaxies which are offset from the main sequence of star-forming galaxies. Since carbon can also be significantly depleted on carbon-rich dust grains, the use of O-based gas tracers might instead be preferred.  The [O{\\sc{i}}]$_{63}$ line has a critical density $n_{\\text{crit,H}}$ $\\sim$ 5 $\\times$ 10$^{5}$ cm$^{-3}$ and upper state energy $E_{\\text{u}}/k$ $\\sim$ 228 K (see Table \\ref{critical}), which makes it an efficient coolant in dense and/or warm photo-dissociation regions (PDRs). Although generally observed to be the second brightest line (after [C{\\sc{ii}}]), the [O{\\sc{i}}]$_{63}$ line is observed to be brighter in galaxies with warm FIR colors and/or high gas densities \\citep{2001ApJ...561..766M,2008ApJS..178..280B,2012A&A...548A..91L}. The applicability of [O{\\sc{i}}]$_{63}$ as a SFR calibrator might, however, be hampered by self-absorption \\citep{1996AAS...189.2102K,1996ApJ...462L..43P}, optical depth effects (more so than [C{\\sc{ii}}]) and the possible excitation of [O{\\sc{i}}]$_{63}$ through shocks \\citep{1989ApJ...342..306H}. In the situation that the gas heating is not longer dominated by the photo-electric effect but has an important contribution from other heating mechanisms (e.g. mechanical heating, soft X-ray heating), the origin of the line emission might differ from the paradigm of warm and/or dense PDRs. A possible origin of [O{\\sc{i}}]$_{63}$ emission different from PDRs can be, in particular, expected in the most metal-poor galaxies characterized by overall low metal content and diminished PAH abundances.  With an energy of the upper state $E_{\\text{u}}/k$  $\\sim$ 163 K, critical density $n_{\\text{crit,e}}$ $\\sim$ 510 cm$^{-3}$ and a high ionization potential of 35.1 eV for O$^{+}$, the [O{\\sc{iii}}]$_{88}$ line originates from diffuse, highly-ionized regions near young O stars. Ionized gas tracers such as [O{\\sc{iii}}]$_{88}$ might gain in importance in low-metallicity environments where PDRs occupy only a limited volume of the ISM judging from their weak PAH emission \\citep{2004A&A...428..409B,2005ApJ...628L..29E,2006ApJ...646..192J,2006A&A...446..877M,2007ApJ...663..866D,2008ApJ...678..804E,2008ApJ...672..214G} and low CO abundance \\citep{1995ApJ...454..293P,1996ApJ...465..738I,1997ApJ...483..200M,2011A&A...531A..19I}. Based on our \\textit{Herschel} observations, we focus on [O{\\sc{iii}}]$_{88}$ as ionized gas tracer. Due to the low critical densities for [O{\\sc{iii}}]$_{88}$ excitation with electrons, other lines (e.g. [O{\\sc{iii}}]$_{52}$ but also optical lines such as [O{\\sc{iii}}] 5007$\\AA$ and H$\\alpha$) will likely dominate the cooling of ionized gas media for intermediate and high gas densities. The brightness of the [O{\\sc{iii}}]$_{88}$ emission line in metal-poor galaxies (Cormier et al. in prep.), however, hints at a SFR tracer with great potential for the high-redshift Universe.  On top of the SFR calibrations for single FIR lines, we try to combine the emission of [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ lines to trace the SFR. The total gas cooling in galaxies scales with the SFR assuming that the ISM is in thermal equilibrium and the total cooling budget balances the gas heating. Therefore, any FIR fine-structure line can be considered a reliable tracer of the SFR if it plays an important role in the cooling of gas that was heated by young stellar photons. Other heating mechanisms unrelated to the UV radiation (e.g. mechanical heating, cosmic ray heating and X-ray heating) and, thus, not directly linked to SFR, might disperse this link.  Ideally, we combine the emission of several cooling lines in the ultraviolet/optical (e.g. Ly $\\alpha$, H$\\alpha$, [O{\\sc{iii}}] 5007$\\AA$) and infrared (e.g. [Ne{\\sc{iii}}] 16\\,$\\mu$m, [S{\\sc{iii}}] 19, 33\\,$\\mu$m, [C{\\sc{ii}}] 158\\,$\\mu$m,  [O{\\sc{i}}] 63, 145\\,$\\mu$m, [N{\\sc{ii}}] 122, 205\\,$\\mu$m, [O{\\sc{iii}}] 52, 88\\,$\\mu$m, [N{\\sc{iii}}] 57\\,$\\mu$m and [Si{\\sc{ii}}] (35\\,$\\mu$m)) wavelength domains to cover the total gas cooling budget in galaxies. Given our focus on the far-infrared coolants, we attempt to obtain a more complete picture of the overall cooling budget by combining the [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ lines.  The \\textit{Herschel} Dwarf Galaxy Survey (DGS) and observations are presented in Section \\ref{data}, together with the acquisition and processing of ancillary data, which will be used as reference SFR calibrators. In Section \\ref{Res.sec}, we take advantage of the high spatial resolution attained for the most nearby galaxies to study the trends and scatter in the spatially resolved relation between the SFR and FIR line emission. The SFR calibrations and trends with the scatter in the SFR-$L_{\\text{line}}$ relations based on global galaxy measurements for the entire DGS sample are presented in Section \\ref{Int.sec}. SFR calibrations for each of the fine-structure lines [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ are derived for a selection of different galaxy populations in Section \\ref{Extend.sec}. In Section \\ref{Conclusions.sec}, we draw together our conclusions. Appendix \\ref{DGScompare} discusses the applicability of different unobscured (Section \\ref{compareSFR}) and obscured (Section \\ref{compareSFRobsc}) indicators as reference SFR tracers for the low-metallicity DGS sample. Appendix \\ref{Compare.sec} presents a comparison between \\textit{Herschel} and ISO spectroscopy. Tables with source information, measurements of reference SFR calibrators and FIR lines are presented in Appendix \\ref{tables} for the DGS sample, the literature sample of galaxies with starburst, composite or active galactic nucleus classifications and high-redshift sources.  \\begin{table*} \\caption{Excitation conditions of the fine-structure lines [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$, with column 2 providing the ionization potential (IP) to create the species. Columns 3, 4 and 5 give the energy of the upper state and critical density for line excitation through collisions with hydrogen atoms and electrons, respectively, calculated for a kinetic temperature $T_{\\text{kin}}$ $=$ 8000K. Column 6 provides the reference to the literature work reporting the excitation conditions of each fine-structure line. Column 7 specifies the ISM phase(s) from which the lines might originate.} \\label{critical} \\centering \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline Line & IP\\tablefootmark{a} & $E_{\\text{u}}/k$ & $n_{\\text{crit,H}}$ & $n_{\\text{crit,e}}$ & Ref\\tablefootmark{b} & Origin? \\\\ \\hline & [eV] & [K] & [cm$^{-3}$] & [cm$^{-3}$] & & \\\\ \\hline $[{\\rm CII}]$ & 11.3 & 91 & 1.6 $\\times$ 10$^{3}$ & 44 & 1 & PDRs, diffuse H{\\sc{i}} clouds, diffuse ionized gas, H{\\sc{ii}} regions  \\\\ $[{\\rm OI}]_{63}$ & - & 228 & 5 $\\times$ 10$^5$ & - & 2 & warm and/or dense PDRs \\\\ $[{\\rm OIII}]_{88}$ & 35.1 & 163 & - & 510 & 3 & low excitation, highly ionized gas \\\\ \\hline \\end{tabular} \\tablefoot{ \\tablefoottext{a}{\\footnotesize The ionization potential refers to the ionization energy of an atom to create the species in Column 1. For example, the ionization potential of 35.1 eV for [O{\\sc{iii}}]$_{88}$ indicates the energy required to remove another electron from O$^{+}$ and, thus, ionize O$^{+}$ in order to create O$^{++}$.}\\\\ \\tablefoottext{b}{\\footnotesize References: (1) \\citet{2012ApJS..203...13G}; (2) \\citet{1985ApJ...291..722T}; (3) \\citet{1999ApJS..123..311A}.}} \\end{table*} ",
        "conclusions": "\\label{Conclusions.sec}  Based on \\textit{Herschel} observations of low-metallicity dwarf galaxies from the Dwarf Galaxy Survey, we have analyzed the applicability of FIR fine-structure lines to reliably trace the star formation activity. More specifically, we investigated whether three of the brightest cooling lines in the DGS sample ([C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$, [O{\\sc{iii}}]$_{88}$) are linked to the star formation rate as probed through a composite SFR tracer (GALEX $FUV$+MIPS 24\\,$\\mu$m). We briefly summarize the results of our analysis: \\begin{itemize} \\item On spatially resolved galaxy scales, the [O{\\sc{iii}}]$_{88}$ line shows the tightest correlation with the SFR (0.25 dex), which provides determination of the SFR with an uncertainty factor of 1.6. Also [O{\\sc{i}}]$_{63}$ is a reasonably good SFR tracer with an uncertainty factor of 1.7 on the SFR estimate. The spatially resolved relation between [C{\\sc{ii}}] and the SFR is heavily dispersed and does not allow us to constrain the SFR within a factor of 2. \\item The dispersion in the SFR calibrations results from the diversity in ISM conditions (i.e. density and ionization state of the gas) for the DGS sample covering a wide range in metallicity, rather than from variations within one single galaxy (on spatially resolved scales). \\item On global galaxy scales, the dispersion in the SFR-$L_{\\text{[CII]}}$ relation (0.38 dex) is again worse compared to the [O{\\sc{i}}]$_{63}$ (0.25 dex) and [O{\\sc{iii}}]$_{88}$ (0.30 dex) lines. The [O{\\sc{i}}]$_{63}$ line is the most reliable overall SFR indicator in galaxies of sub-solar metallicity with an uncertainty factor of 1.8 on the SFR estimate, while the SFR derived from [O{\\sc{iii}}]$_{88}$ is uncertain by a factor of 2. The [C{\\sc{ii}}] line is not considered a reliable SFR tracer in galaxies of low metal abundance. \\item The scatter in the SFR-$L_{\\text{[CII]}}$ relation increases towards low metallicities, warm dust temperatures and large filling factors of diffuse, highly ionized gas. Due to the porosity of the ISM and the exposure to hard radiation fields, an increased number of ionizing photons is capable of ionizing gas at large distances from the star-forming regions, which favors line cooling through ionized gas tracers such as [O{\\sc{iii}}]$_{88}$. The photo-electric efficiency might, furthermore, reduce in low-metallicity environments due to grain charging and/or increased photon escape fractions. \\item On spatially resolved scales, we can reduce the scatter in the SFR calibration by combining the emission from multiple FIR lines. Ideally, we want to probe the emission from all cooling lines that constitute the total gas cooling budget. \\end{itemize}  Based on the assembly of literature data, we furthermore analyze the applicability of fine-structure lines [C{\\sc{ii}}], [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ to probe the SFR (as traced by the TIR luminosity) in H{\\sc{ii}}/starburst galaxies, AGNs, ULIRGs and high-redshift objects: \\begin{itemize} \\item The [C{\\sc{ii}}] and [O{\\sc{i}}]$_{63}$ lines are considered to be the most reliable SFR tracers to recover the star formation activity in starburst galaxies with uncertainty factors of 1.9 and 1.6, respectively. The [O{\\sc{iii}}] line, on the other hand, is weak and the SFR calibration could not be well constrained due to the low number of \\textit{Herschel} [O{\\sc{}iii}]$_{88}$ detections in starbursts. \\item All three FIR lines can recover the SFR from composite or AGN sources within an uncertainty of factor $\\sim$ 2.3. The increased scatter in the SFR calibrations for AGNs (as compared to starbursts) might result from a possible AGN contribution to the total infrared luminosity (used to derive the SFR). Alternatively, some AGNs might show line deficits similar to ULIRGs. \\item ULIRGs are offset from the SFR calibrations for starbursts and AGNs due to line deficits relative to their total-infrared luminosity and, therefore, require separate SFR calibrations. The star formation rate in ULIRGs is preferentially traced through [C{\\sc{ii}}] and [O{\\sc{i}}]$_{63}$ line emission, providing SFR estimates with uncertainties of factor $\\sim$ 2, while the SFR([O{\\sc{iii}}]) estimate is uncertain by a factor of 2.5. \\item At high-redshift, we can only reliably determine a SFR calibration for the [C{\\sc{ii}}] line (with an uncertainty factor of $\\sim$ 2.5 on the SFR estimate), due to the low number of observations for the other lines. The relatively few detections of [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ appear to be bright at high-redshift, suggesting that the [O{\\sc{i}}]$_{63}$ and [O{\\sc{iii}}]$_{88}$ lines are also potentially powerful tracers of the SFR at high redshift, but more detections are mandatory to acquire conclusive evidence. \\end{itemize}"
    },
    "1402/1402.6593_arXiv.txt": {
        "abstract": "Magnetic Rayleigh-Taylor (MRT) instabilities may play a relevant role in many astrophysical problems. In this work the effect of magnetic shear on the growth rate of the MRT instability is investigated. The eigenmodes of an interface and a slab model under the presence of gravity are analytically calculated  assuming that the orientation of the magnetic field changes in the equilibrium, i.e., there is magnetic shear. We solve the linearised magnetohydrodynamic (MHD) equations in the incompressible regime.We find that the growth rate is bounded under the presence of magnetic shear. We have derived simple analytical expressions for the maximum growth rate, corresponding to the most unstable mode of the system. These expressions provide the explicit dependence of the growth rate on the various equilibrium parameters. For small angles the growth time is linearly proportional to the shear angle, and in this regime the single interface problem and the slab problem tend to the same result. On the contrary, in the limit of large angles and for the interface problem the growth time is essentially independent of the shear angle. In this regime we have also been able to calculate an approximate expression for the growth time for the slab configuration. Magnetic shear can have a strong effect on the growth rates of the instability. As an application of the results found in this paper we have indirectly determined the shear angle in solar prominence threads using their lifetimes and the estimation of the Alfv\\'en speed of the structure. ",
        "introduction": "\\label{sec:intro}  The magnetic Rayleigh-Taylor (MRT) instability is important in many astrophysical systems. Some examples are buoyant magnetised bubbles identified in clusters of galaxies, see \\citet[][]{robinson2004,jones2005} for studies in 2D, and \\citet[][]{oneil2009} for 3D configurations. MRT instabilities also manifest themselves in shells of young supernova remnants, this has been investigated by  \\citet[][]{jun1995} in 2 and 3D Cartesian configurations and by \\citet[][]{jun1996} in 3D using spherical coordinates. \\citet{bucc2004} have numerically investigated the development of the MRT instability at the interface between an expanding pulsar wind nebula and its surrounding supernova remnant. \\citet{stonegardiner2007} studied the behaviour of magnetic Rayleigh-Taylor instability in three dimensions with special focus on the structure and dynamics of the nonlinear evolution of the system. They analysed various configurations including the situation in which magnetic fields  change direction at the interface between the two fluids.  \\citet{stonegardiner2007} used the MRT instability to explain the structure of the optical filaments observed in the Crab nebula.  In laboratory plasmas the possible stabilising effect by a force-free magnetic field has been studied in the past by many authors \\citep[see for example][]{goedbloed1971a,goedbloed1971b,goedbloed1971c,GOEPOE2004} using the single interface problem and the slab problem and applying vacuum conditions at some of the boundaries. \\citet{yang2011} have studied the magnetic field transition layer effects on the MRT instability with continuous magnetic field and density profiles and have found that the linear growth rate of the MRT instability increases with the thickness of the magnetic field transition layer, especially for the case of small thickness. Recently, \\citet{zhangetal2012} have used the ideal MHD model to study the effect of magnetic shear in a finite slab representing a magnetic liner, which is a device used in experiments with fusion plasmas. These authors have found that magnetic shear reduces the MRT growth rate in general.  The emergence of magnetic flux from the solar interior and the formation of flux tubes is another example where MRT instabilities are relevant. For example, \\citet[][]{isobe2005,isobe2006} proposed that the MRT instability is a possible cause of the filamentary structure in mass and current density in the emerging flux regions. In the solar atmosphere, \\citet{ryuetal10} suggested that several dynamic processes taking place in prominences are most probably related to magnetic Rayleigh-Taylor instabilities. Along this line of work, \\citet{hillier2011,hillier2012a,hillier2012b} have performed three-dimensional magnetohydrodynamic simulations to investigate the nonlinear evolution of the Kippenhahn-Shl\\\"uter prominence model to the MRT instability.  The fine structure of solar prominences reveals the presence of magnetic threads. These structures are quite thin,  of the order of $100\\,\\rm km$, aligned with the magnetic field and, in many cases, they seem to lie horizontally with respect to the photosphere \\citep[see][for recent results about the formation of these structures]{devore2012,devore2013}. \\citet{terretal2012} have considered the possible link between magnetic Rayleigh-Taylor instabilities and the short thread lifetimes. In that work a slab model permeated by a horizontal magnetic field was considered. The growth rates of the unstable modes and the thresholds for stability were determined analytically. In the present paper we extend the study to the situation with a sheared magnetic field in which the magnetic field changes its direction at the interfaces of the plasma slab. To understand the results in the slab model we describe first the effect of shear at a single plasma interface. Magnetic shear introduces changes in the growth rates of the unstable modes that might be relevant regarding the lifetime of threads. In this work we analytically calculate these growth rates and perform a detailed analysis of their dependence on the equilibrium parameters. ",
        "conclusions": "\\label{sec:sum}  In the present paper we have extended the study by \\citet{terretal2012} to the models of an interface and a slab with magnetic shear, and have focused on the MRT instability. The fact that the magnetic field changes its direction introduces a bounded growth rate of the instability. This is different from the models without shear where the growth rate is unbounded. Because of this \\citet{terretal2012} concentrated on a particular wave number in the perpendicular direction equal to $1/a$\\/, $a$ being the half-thickness of the slab. Here we have focused on the maximum growth rate, representing the most unstable mode of the system, and have found analytical expressions in various limiting cases. For small angles of the shear the growth time is linearly proportional to the shear angle $\\alpha$\\/. This applies to the single interface as well as to the slab problem. On the contrary, for large angles the growth time depends only weakly on $\\alpha$ in the interface problem. In this limit, we have also been able to calculate an approximate expression for the growth time in the slab problem.  We have shown, using a simple example, how it is possible to estimate the shear angle in threads belonging to active region prominences from the combination of observations and the theoretical results presented in this paper. Using this method we have found that the observations of oscillating threads of \\citet{okamoto07} are compatible with small shear angles (around $13^{\\circ}$). This indirect method of inferring some of the equilibrium properties of threads can be potentially used as a seismological tool.  Several simplifications have been done in the models considered in this work. Both the interface and slab configurations are unbounded in the perpendicular direction ($y$\\/-direction) and we have assumed that there is dense material along the full length of the magnetic tube. However, in reality, threads only represent a small part of the full magnetic tube. Compressibility and partial ionisation \\citep[see][]{diaz2012} have been also ignored. These basic assumptions in the models have enabled us to derive analytical expressions for the growth times. Nevertheless, the need for using improved models is obvious, and the study of the nonlinear evolution of the system is very relevant to asses the role of the MRT instability in the fast disappearance of threads. \\vspace*{5mm}  \\noindent {\\bf Acknowledgements}. A part of this work was carried out when MSR was a guest of Departament de F\\'{\\i}sica of Universitat de les Illes Balears. He acknowledges the financial support received from the Universitat de les Illes Balears and the warm hospitality of the Departament. He also acknowledges the support by the STFC grant. J.T. acknowledges support from the Spanish Ministerio de Educaci\\'on y Ciencia through a Ram\\'on y Cajal grant and financial support from MICINN/MINECO and FEDER Funds through grant AYA2011-22846 Funding from CAIB through the Grups Competitius scheme and FEDER Funds is also acknowledged."
    },
    "1402/1402.7076_arXiv.txt": {
        "abstract": "The slope of the star formation rate/stellar mass relation (the SFR ``Main Sequence\"; $\\sfr$--$\\Mstel$) is not quite unity: specific star formation rates ($\\sfr/\\Mstel$) are weakly-but-significantly anti-correlated with $\\Mstel$. Here we demonstrate that this trend may simply reflect the well-known increase in bulge mass-fractions -- portions of a galaxy {\\it not} forming stars -- with $\\Mstel$.  Using a large set of bulge/disk decompositions and SFR estimates derived from the Sloan Digital Sky Survey, we show that re-normalizing SFR by {\\it disk} stellar mass ($\\ssfrd\\equiv\\sfr/\\Mdisk$) reduces the $\\Mstel$-dependence of SF efficiency by $\\sim0.25$ dex per dex, erasing it entirely in some subsamples. Quantitatively, we find $\\log\\ssfrd$--$\\log\\Mstel$ to have a slope $\\betare\\in[-0.20,0.00]\\pm0.02$ (depending on $\\sfr$ estimator and Main Sequence definition) for star-forming galaxies with $\\Mstel\\geq10^{10}\\Msun$ and bulge mass-fractions $B/T\\lesssim0.6$, generally consistent with a pure-disk control sample ($\\betacon=\\betaconval$). That $\\langle\\sfr/\\Mdisk\\rangle$ is (largely) independent of host mass for star-forming disks has strong implications for aspects of galaxy evolution inferred from any $\\sfr$--$\\Mstel$ relation, including: manifestations of ``mass quenching\" (bulge growth), factors shaping the star-forming stellar mass function (uniform $d\\log\\Mstel/dt$ for low-mass, disk-dominated galaxies), and diversity in star formation histories (dispersion in $\\sfr(\\Mstel,t)$).  Our results emphasize the need to treat galaxies as composite systems -- not integrated masses -- in observational and theoretical work. ",
        "introduction": "\\label{sec:intro}  The observation of a correlation between galaxy star formation rates (SFRs) and stellar masses ($\\Mstel$) has generated considerable interest. Seen from $z = 0$ to $z > 2$ \\citep[e.g.,][and references therein]{Brinchmann04, Daddi07, Wuyts11, Guo13}, this SFR ``Main Sequence\" (\\MS) may encode fundamental information about galaxy evolution.  Uncontroversial is the fact that the \\MS\\ has fallen monotonically since at least $z\\sim2$ \\citep[e.g.,][]{Noeske07, Rodighiero10, Whitaker12}.  Observed at all $\\Mstel\\gtrsim10^{10}\\Msun$, this phenomenon must contribute significantly to the precipitous decline in cosmic star formation seen over the same epoch \\citep[e.g.,][]{Lilly96, Madau96, Cucciati12}.  However, while its gross evolution is increasingly well-understood, the slope and dispersion of the \\MS\\ remain uncertain. Such uncertainty arises (at least) from dependencies on SFR indicators \\citep[e.g.,][Figure 4]{Pannella09}, the definition of ``star forming\" \\citep[e.g.,][\\S7.5]{Salim07}, and a lack of high-redshift data at moderate-to-low $\\Mstel$ \\citep[e.g.,][Figure 1]{Whitaker12}.  Despite these issues, if its evolution reflects that of individual systems, the slope and dispersion of the \\MS, their time-dependence, and their interpretation have deep implications for pictures of galaxy growth. With the \\MS\\ broadly reproducible in cosmological simulations \\citep[][]{Keres05, Neistein08, Lagos11, Hopkins13} and actively employed as a basis/constraint for evolutionary models \\citep[][]{PengLilly10, Leitner12, Behroozi13}, understanding such details is increasingly important.  Here we reinterpret the slope of the \\MS.  \\begin{deluxetable*}{lccl} \\tablecolumns{4} \\tablecaption{Quantities} \\tablehead{ \\colhead{Quantity} & \\colhead{Unit} & \\colhead{Source\\tablenotemark{a}} & \\colhead{Definition} } \\startdata $r_{\\rm disk}$\t\t& SDSS mag\\tablenotemark{b}\t& 1 & Disk absolute $r$ magnitude $k$-corrected to $z=0.1$\\\\ $(g-r)_{(\\rm disk)}$\t& SDSS mag\\tablenotemark{b}\t& 1 & (Disk) rest-frame color $k$-corrected to $z=0.1$\\\\ $z$\t\t\t\t& $\\cdots$\t\t\t\t\t& 1 & Galaxy redshift\\\\ $\\ns$\t\t\t& $\\cdots$\t\t\t\t\t& 1 & Global $r$-band \\sersic index\\\\ $b/a$\t\t\t& $\\cdots$\t\t\t\t\t& 1 & Global $r$-band axis ratio (1 - ellipicity)\\\\ $M_{\\ast(,\\,\\rm disk)}$& $\\Msun$\t\t\t\t& 2(3) & (Disk) stellar mass\\\\ SFR \t\t\t\t& $\\Msun\\, {\\rm yr^{-1}}$ \t\t& 2 & Aperture-corrected star formation rate (median of PDF)\\\\ sSFR \t\t\t\t& ${\\rm yr^{-1}}$ \t\t\t& 2 & Galaxy specific star formation rate ($\\sfr/\\Mstel$) \\\\ $\\ssfrd$\t\t\t& ${\\rm yr^{-1}}$ \t\t\t& 3 & Disk-mass-normalized star formation rate ($\\sfr/\\Mdisk$) \\enddata \\tablenotetext{a}{1--S11; 2--\\bof/\\bos; 3--Derived.} \\tablenotetext{b}{AB system offsets are $<0.01$ mag.} \\label{tbl:params} \\end{deluxetable*}  The \\MS\\ is conveniently recast in terms of galaxies' {\\it specific} star formation rates -- $\\ssfr\\equiv\\sfr/\\Mstel$ -- or fractional mass-growth per unit time. If constant, $\\ssfr$ is the (inverse) $\\Mstel$ $e$-folding timescale.  The $\\Mstel$-dependence of sSFR -- the departure of the \\MS\\ slope from unity -- contains information about the ``efficiency\" of SF across the galaxy mass spectrum.\\footnote[5]{More direct definitions of ``SF efficiency\" relate $\\sfr$ to a gas mass, but $\\ssfr$ is {\\it an} efficiency metric.}  Typically, it is parametrized by the power-law index: \\begin{equation} \\beta\\equiv\\frac{d\\log\\ssfr}{d\\log\\Mstel}. \\end{equation}  If all galaxies formed stars with equal efficiency, $\\beta$ would be identically zero.  Observationally, $\\beta$ appears {\\it close} to zero, permitting convenient approximations in evolutionary models \\citep[e.g.,][]{PengLilly10}; $\\ssfr(t)$ is {\\it nearly} independent of mass, so the entire star-forming population is {\\it nearly} describable by a single number (absent significant dispersion at fixed $\\Mstel$; see Section \\ref{sec:scatter}).  Yet, $\\beta$ is {\\it not} zero. Many studies using SFR indicators from the UV through the radio have concluded that, above $10^{10}\\Msun$, $-0.6\\lesssim\\beta\\lesssim-0.1$ for $z\\lesssim2$ \\citep[][but cf.\\ \\citealt{Pannella09}]{Brinchmann04, Salim07, Karim11, Whitaker12}. \\citet[][]{PengLilly10} and \\citet[][]{Whitaker12} find $\\beta\\simeq0$ for blue galaxies (see Section \\ref{sec:implications} below), but that $\\beta$ is significantly negative for the global star-forming population seems secure.  The implication of $\\beta<0$ is that low-mass galaxies grow (logarithmically) faster than higher-mass contemporaries.  Interesting on its own, this fact is important also because $\\beta$ informs two other key questions: Why has the shape of the star-forming stellar mass function remained unchanged since $z\\sim2$ \\citep[e.g.,][]{Ilbert10, Tomczak14}? What {\\it stops} star formation?  Setting aside the mass function for now (see Section \\ref{sec:implications} and the extensive treatment of \\citealt{PengLilly10}) the question of what stops SF in galaxies nicely illustrates $\\beta$'s influence on \\MS-based evolutionary models.  In the $\\beta\\rightarrow0$ limit, galaxy evolution is binary: systems are either star-forming -- growing in lock-step with all other such objects -- or not.  An implication is that mechanisms taking galaxies from the first population into the second act quickly and operate across all $\\Mstel$.  Conversely, if $\\beta$ is substantially negative (as is likely), galaxy evolution is more nuanced.  A galaxy's global SF efficiency changes as it grows, gradually falling to negligible levels with time. ``Quenching\" is thus a mix of processes pulling systems vertically off the \\MS\\ and lowering $\\ssfr$s as they move along it ($\\beta$ reflects the latter).  Many mechanisms have been proposed that implicitly or explicitly account for mass-dependent $\\ssfr$s, including virial-heating of the circumgalactic medium by dark matter halos (inducing ``hot-mode\" accretion) and AGN activity \\citep[e.g.,][]{Dekel04, Keres05, Croton08, vandeVoort11b}.  Such processes may be at work, but they are not directly coupled to the observables in $\\ssfr$--$\\Mstel$, so hypotheses are complicated by uncertainties in linking these phenomena.  Indeed, observationally, there is a deeper concern. $\\ssfr=\\sfr/\\Mstel$ (hence $\\beta$) is biased, {\\it prima facie}, as a description of SF as the numerator has essentially nothing to do with a significant part of the denominator -- the bulge. Given the well-known correlation of bulge mass-fractions, $B/T$, with $\\Mstel$, $\\beta<0$ is expected simply because ever smaller portions of a galaxy participate in SF, independent of the nature of the SF itself.  If $\\ssfr$--$\\Mstel$ is to add meaningfully to our knowledge of galaxy evolution, at a minimum, the extent to which $\\beta$ reflects changes in the quality of SF ({\\it how}) versus the proportion of a galaxy contributing to it ({\\it where}) must be understood.  Large spectrophotometric surveys -- such as the Sloan Digital Sky Survey \\citep[SDSS --][]{York00} -- enable this.  Below, we demonstrate the importance of recognizing {\\it where} SF occurs, showing that most-to-all of $\\beta$ can be erased simply by redefining ``$\\ssfr$\" using the mass in galactic disks. ",
        "conclusions": "Re-normalizing $\\sfr$ by {\\it disk} stellar mass, $\\Mdisk$ can account for $\\sim0.25$ dex of declining $\\ssfr$ per decade $\\Mstel$, essentially removing the dependence of SF efficiency on galaxy mass for star-forming systems with blue disks (if not all star-forming galaxies).  Besides suggesting a key part of ``mass-quenching\" is ``bulge-building\" -- distinct from processes affecting SF in disks -- our findings ease tension between the \\MS\\ and the evolution of the stellar mass function, and reinforce two important points: \\begin{itemize} \\item ``Understanding galaxy evolution demands the routine bulge--disk decomposition of the giant galaxy population at all redshifts,\" \\citep{Allen06}; \\item Dispersion in $\\sfr(\\Mstel)$ likely reflects real diversity in SFHs and should not be ignored. \\end{itemize}  Upcoming IFU surveys (e.g., MaNGA; \\url{www.sdss3.org/future/manga.php}) may constrain intrinsic spreads in $\\sfr(M_{\\ast(,\\rm disk)}$) and thus mechanisms shaping SFHs. Regardless, $\\sfr/\\Mdisk$--$\\Mstel$ and $B/T$--$\\Mstel$ should serve as benchmarks for future theoretical models of galaxy evolution."
    },
    "1402/1402.6995_arXiv.txt": {
        "abstract": "We consider super-critical accretion with angular momentum onto stellar-mass black holes as a possible mechanism for growing billion-solar-mass holes from light seeds at early times. We use the radiatively-inefficient ``slim disk\" solution -- advective, optically thick flows that generalize the standard geometrically thin disk model -- to show how mildly super-Eddington intermittent accretion may significantly ease the problem of assembling the first massive black holes when the universe was less than 0.8 Gyr old. Because of the low radiative efficiencies of slim disks around non-rotating as well as rapidly rotating holes, the mass $e$-folding timescale in this regime is nearly independent of the spin parameter. The conditions that may lead to super-critical growth in the early universe are briefly discussed. ",
        "introduction": "\\label{intro}  The most distant quasar discovered to date, ULAS J1120+0641 at a redshift $z=7.084$, is believed to host a black hole with a mass of $2.0^{+1.5}_{-0.7}\\times 10^9\\,\\msun$ that is shining 0.78 Gyr after the big bAng \\citep{Mortlock11}. This object, together with the handful of bright {Sloan Digital Sky Survey} (SDSS) quasars at redshift $z\\gta 6$ \\citep{Fan06}, sets some of the tightest constraints on models for the formation and growth of massive black holes (MBHs) at early epochs. The challenge provided by the existence of billion-solar-mass black holes at the end of the reionization epoch is easily described (see, e.g., \\citealt{Haiman13} for a review). If MBHs are assembled by the accretion of gas onto less massive ``seed\" holes at the rate $\\dot m$, and if in the process a fraction $\\epsilon$ of the rest-mass energy of the infalling material is released as radiation, then the growth of the hole's mass $M$ is regulated by the equation \\begin{equation} {dM\\over dt}=(1-\\epsilon)\\dot m=\\left({1-\\epsilon\\over \\epsilon}\\right)\\left({L\\over L_E}\\right){M\\over t_E}. \\end{equation} Here $L$ is the radiated luminosity, $\\epsilon \\equiv L/\\dot m c^2$, $L_E \\equiv 4\\pi G M \\mu_e m_pc/\\sigma_T$ is the Eddington limit when the continuum radiation force balances gravity, $\\sigma_T$ is the Thomson scattering cross-section, $\\mu_e$ is the mean molecular weight per electron, $t_E\\equiv Mc^2/L_E=0.44\\mu_e^{-1}$ Gyr is the Eddington timescale, and all other symbols have their usual meaning. The characteristic $e$-folding timescale $t_{\\rm acc}$ for mass growth is then \\begin{equation} t_{\\rm acc}=\\left({\\epsilon\\over 1-\\epsilon}\\right)\\left({L_E\\over L}\\right)t_E=(4.3\\times 10^7~{\\rm yr})\\,\\left({L_E\\over L}\\right), \\label{eq:tacc} \\end{equation} where the last equality assumes $\\mu_e=1.15$ (valid for primordial gas) and a radiative efficiency of $\\epsilon=0.1$.  In a concordance cosmology with $\\Omega_M=0.27$, $\\Omega_\\Lambda=0.73$, and $H_0=70\\,\\kmsmpc$, the time elapsed between $z=20$ and $z=7$ is 0.6 Gyr, corresponding to 14 $e$-foldings of Eddington-limited accretion ($L=L_E$) and a mass amplification factor of $10^{6}$.  The growth of $2\\times 10^9\\,\\msun$ MBHs at the Eddington rate from light black hole seeds of mass $M_0=100\\,\\msun$ requires $\\ln(2\\times 10^9/100)=17$ $e$-foldings at $L=L_E$. Therefore, if the first seeds were $\\sim 100\\,\\msun$ remnants of the first generation of massive stars \\citep[e.g.,][]{Madau01,Haiman01,Heger03,Volonteri03}, these could grow into billion-solar-mass holes by $z\\sim 7$ only if all the following conditions were fulfilled: (1) seeds were present early on, at $z\\gta 20$; (2) gas accretion continued more or less uninterrupted at the Eddington rate for $\\gta 0.6$ Gyr; and (3) $\\epsilon<0.1$ \\citep{Tanaka09}. The second condition is hard to satisfy in the shallow potential wells of low-mass dark matter halos, as feedback effects resulting from the accretion process itself are expected to dramatically affect gas inflow and may result in sub-Eddington rates and negligible mass growth \\citep{Johnson07,Pelupessy07,Alvarez09,Milos09}. The third condition requires radiative efficiencies that are below those expected for thin disk accretion onto rapidly spinning Kerr black holes ($\\epsilon \\simeq 0.3-0.4$, \\citealt{Thorne74,Shapiro05}), and approach the value, $\\epsilon=0.057$, characteristic of the Schwarzschild non-rotating solution.  Over the last decade, a number of alternatives to the above picture have been proposed. If stellar seeds were present in large numbers at high redshifts, coalescing black hole binaries brought together by successive galaxy mergers may, in principle, help mass build-up and generate mass amplification factors as high as $10^4$ \\citep{Yoo04}. More massive seeds, with $M_0\\sim 10^4-10^5\\,\\msun$, may form through the ``direct collapse\" of low angular momentum gas at high redshift \\citep[e.g.,][]{Loeb94,Bromm03,Koushiappas04,Lodato06,Spaans06,Regan09,Mayer10}, likely via the intermediate stage of supermassive stars \\citep{Begelman10}, and therefore ``jump start\" the whole process.  Questions remain about the idealized conditions needed in these models to avoid fragmentation, dissipate angular momentum, and drive gas towards the center of protogalaxies at extremely high rates.  In this {Letter} we discuss {\\it super-critical (i.e., super-Eddington) accretion with angular momentum onto stellar-mass seeds} as a possible mechanism for bypassing some of the above difficulties. Evidence for near-Eddington or super-Eddington flows has been accumulating in recent years. Super-critical accretion onto stellar-mass black holes has been invoked to explain the nature of the ultraluminous X-ray sources \\citep[e.g.,][]{Gladstone09,Middleton13}. A study of a large sample of active galactic nuclei (AGNs) suggests that many of them emit considerably more energy and have higher $L/L_E$ ratios than previously assumed \\citep{Netzer14}. \\citet{Kormendy13} have recently argued that the normalization of the local black hole scaling relations should be increased by a factor of five to $M=0.5\\%\\,M_{\\rm bulge}$. This increases the local mass density in black holes by the same factor, decreases the required mean radiative efficiency to 1\\%-2\\%, and may be evidence for radiatively-inefficient super-Eddington accretion \\citep[e.g.,][]{Soltan82,Novak13}. At high redshifts, the very soft X-ray spectrum of ULAS J1120+0641 appears to suggest that this quasar is accreting at super-critical rates \\citep{Page13}. On the theoretical side, it is known that the standard, radiatively-efficient thin disk solution \\citep{Shakura76} can no longer be applied when mass is supplied to a black hole at high rates. In this regime, viscosity-generated heat does not have sufficient time to be radiated away, and is instead advected into the hole. The shorter mass $e$-folding timescales and the decreased radiative efficiencies that characterize these flows make them ideal for feeding and growing MBHs out of stellar-mass seeds at early times. ",
        "conclusions": "The assembly of pregalactic MBHs following a phase of super-critical quasi-spherical accretion in metal-free halos has been explored by \\citet{VolonteriRees05}. Highly super-critical rates at early times have been recently advocated by \\citet{Volonteri14}. Here, we have extended and updated these works by focusing on the ``slim disk\" solution --  advective, optically thick disks that generalize the standard model of radiatively efficient thin flows to moderately super-Eddington accretion rates. Under the assumption that a gas supply rate of a few $\\times\\, \\dot m_E\\approx 0.01\\,M_5\\,\\sfr$ (where $M_5$ is the mass of the hole in units of $10^5\\,\\msun$) is indeed able to reach the central MBH for a period of 20 Myr or so, we have shown how a few episodes of super-critical accretion may turn a handful of light seed holes into the population of rare bright quasars observed at $z\\gta 6$ by the SDSS. Such high rates may both be determined by local physics within the accretion flow as well as by the large-scale cosmological environment in which MBHs and their hosts are growing. An example of small-scale physics is the presence of density inhomogeneities induced by radiative-hydrodynamic instabilities that reduce the effective opacity and allow super-Eddington fluxes from slim, porous disks \\citep{Dotan11}. On large scales, gas-rich major mergers between massive protogalaxies have been shown to drive large gas flows towards the center of the merger remnant, owing to continuous losses of angular momentum by torques and shocks over a wide range of spatial scale \\citep[e.g.,][]{Mayer10}. Contrary to the formation of MBH seeds by rapid ``direct collapse\", this gas supply is not limited by the Kelvin-Helmholtz contraction timescale of the central supermassive proto-star \\citep[see, e.g.,][]{Shang10}, but can add to the central black hole mass over the longer accretion timescale of Equation (\\ref{eq:tsacc}). And while these high-accretion episodes may generate AGN outbursts, drive hot, high-velocity outflows from the host galaxy, and occasionally clear the nuclear region of most of its gas, additional gas inflows may quickly rebuild the nuclear reservoir as AGN feedback appears to couple only weakly to a galaxy disk \\citep{Gabor14}.  On intermediate scales, models of circumnuclear disks supported by the turbulent pressure caused by supernova explosions predict high mass supply rates to the central 1 pc lasting for $10^8$ yr \\citep{Kawakatu08}. Even in the absence of star formation, thick, turbulent pressure-supported disks are often seen in simulations of early protogalaxies \\citep{Wise08}. Viscous torques generated by the dissipation of supersonic turbulent motions generate a mass accretions rate \\citep{Pringle81} \\begin{equation} \\dot m=2\\pi \\nu \\Sigma_0 \\, \\vert \\frac{d\\ln \\Omega}{d\\ln r} \\vert, \\label{eq:feeding} \\end{equation} where $\\Sigma_0$ is the central surface density of the disk, $\\nu \\simeq \\sigma h$ is the viscous parameter \\citep{Wada02}, $\\sigma$ is the gas turbulent velocity, $h$ is the disk scale height, $h/r\\sim \\sigma/V_{\\rm rot}$, and $V_{\\rm rot}$ is the disk rotational velocity. Observations show that, at a fixed stellar mass, galaxies systematically increase in disordered motions and decrease in rotation velocity with increasing redshift and decreasing potential \\citep{Kassin12}. At $z\\sim 2$, the ratio of the rotational to random velocities ranges between 1 and 6, quite in contrast to present-day disk galaxies where $V_{\\rm rot}/\\sigma\\sim 10-20$ \\citep{Forster09}.  Assuming $V_{\\rm rot}/\\sigma\\sim 1$ in the early universe, writing the disk mass as $M_d=2\\pi \\Sigma_0 R_d^2$ where $R_d$ is the disk scale length, and relating it to $V_{\\rm rot}$ using the baryonic Tully-Fisher relation, $M_d={\\cal A}\\,V_{\\rm rot}^4$ \\citep{McGaugh05}, Equation (\\ref{eq:feeding}) yields an order-of-magnitude estimate of the accretion rate expected in this ``turbulent feeding\" model, \\begin{align} \\dot m \\sim {\\cal A} V_{\\rm rot}^5 {r\\over R_d^2} & \\sim (1\\,\\sfr)\\,\\left({V_{\\rm rot}/30\\,\\kms}\\right)^5 \\times \\nonumber \\\\ & \\times \\left({r\\over 0.1\\,R_d}\\right)\\left({{\\rm 0.1\\,kpc}\\over R_d}\\right), \\label{eq:supply} \\end{align} where ${\\cal A}=50\\,\\msun {\\rm km^{-4}\\,s^4}$ \\citep{McGaugh05} and we have dropped numerical factors of order unity.  The estimate above shows that large mass supply rates in the dense environments of high redshift, dispersion-dominated, massive protogalaxies are theoretically plausible. It is consistent with the mass accretion rates measured at 100 pc from the center of simulated atomic cooling halos at $z>10$ \\citep{Prieto13}, and is comparable to the mass accretion rates, $\\sim c_s^3/G\\propto T^{3/2}/G$ expected in self-gravitating, collapsing $T\\sim 10^4\\,$K gas clouds. Rapid MBH growth may occur in rare special environments where gas can flow toward the center at super-Eddington rates relatively unaffected by feedback processes. Indeed it would seem fortuitous if the mass supply rate was always regulated exactly at $\\dot m_E$ during the early growth of MBHs, as it is commonly assumed. Equally contrived, of course, are gas fueling rates that remain at $(3-4)\\times \\dot m_E$ for 20-50 Myr while the central MBHs are growing by many orders of magnitude in mass. Perhaps a more plausible scenario is one where $\\dot m\\gg \\dot m_E$ early on (when black holes are still of intermediate mass), but only a fraction of this external mass supply reaches the horizon, and where $\\dot m$ ``dwindles\" to a few times Eddington during the last few $e$-foldings, when the MBH is already above $10^8\\,\\msun$."
    },
    "1402/1402.6671_arXiv.txt": {
        "abstract": "The eXciting Dark Matter (XDM) model was proposed as a mechanism to efficiently convert the kinetic energy (in sufficiently hot environments) of dark matter into e+e- pairs. The standard scenario invokes a doublet of nearly degenerate DM states, and a dark force to mediate a large upscattering cross section between the two. For heavy ($\\sim \\tev$) DM, the kinetic energy of WIMPs in large (galaxy-sized or larger) halos is capable of producing low-energy positrons. For lighter dark matter, this is kinematically impossible, and the unique observable signature becomes an X-ray line, arising from $\\chi \\chi \\rightarrow \\chi^* \\chi^*$, followed by $\\chi^* \\rightarrow \\chi \\gamma$. This variant of XDM is distinctive from other DM X-ray scenarios in that it tends to be most present in more massive, hotter environments, such as clusters, rather than nearby dwarfs, and has different dependencies from decaying models. We find that it is capable of explaining the recently reported X-ray line at 3.56 keV. For very long lifetimes of the excited state, primordial decays can explain the signal without the presence of upscattering. Thermal models freeze-out as in the normal XDM setup, via annihilations to the light boson $\\phi$. For suitable masses the annihilation $\\chi \\chi \\rightarrow \\phi \\phi$ followed by $\\phi \\rightarrow SM$ can explain the reported gamma-ray signature from the galactic center. Direct detection is discussed, including the possibility of explaining DAMA via the ``Luminous'' dark matter approach. Quite generally, the proximity of the 3.56 keV line to the energy scale of DAMA motivates a reexamination of electromagnetic explanations. Other signals, including lepton jets and the modification of cores of dwarf galaxies are also considered. ",
        "introduction": "The evidence for an additional non-baryonic component of the universe has long since been overwhelming. With data from scales as small as dwarf galaxies to the large-scale properties of the universe, from late time evolution of the universe to the properties of the CMB, the Cold Dark Matter (CDM) paradigm has passed numerous tests.  However, although its existence has been well established, its nature remains a total mystery. While a vanilla dark matter particle - devoid of interesting dynamics - remains a viable possibility, dark matter with nontrivial dynamics has been increasingly considered. Attempts to explain direct detection anomalies, indirect detection anomalies, and anomalies in the properties of dwarf galaxy have all relied on invoking a more complicated dark sector, often with multiple dark states and new, dark forces.  One such scenario was the ``eXciting Dark Matter'' (XDM) proposal \\cite{Finkbeiner:2007kk}. In it, dark matter $\\chi$ is assumed to have an excited state $\\chi^*$, and to interact with itself via a new force with mediator $\\phi$. Scattering via $\\phi$ mediates the excitation $\\chi\\rightarrow \\chi^*$, and the decay $\\chi^* \\rightarrow \\chi + SM$ produces some standard model state, SM.  The idea of XDM was to produce observable signals from the dark matter's {\\em kinetic} energy rather than its mass, using the light mediator to produce detectably large scattering cross sections.\\footnote{A related scenario with excitations into charged excited states $\\chi^\\pm$ was considered in \\cite{Pospelov:2007xh}.}  For splittings $\\delta \\equiv m_{\\chi^*}-m_\\chi \\ge 2 m_e$, the decay $\\chi^*\\rightarrow \\chi e^+ e^-$ can be mediated either by kinetic mixing (for a vector mediator) or Higgs portal mixing (for a scalar mediator), and can produce low energy $e^+ e^-$ pairs. Such a scenario might solve one of the oldest mysteries in high-energy astrophysics: the surprising strength of the 511 keV line and associated positronium continuum in the inner galaxy \\cite{Knodlseder:2003sv,Churazov:2004as,Knodlseder:2005yq}.  In this work, we explore the $\\delta \\le 2 m_e$ case, which can only lead to photon or neutrino final states.  The decay $\\chi^* \\rightarrow \\chi+ \\gamma$ could produce an X-ray line, and in the following we will consider the detectability of such a signal. ",
        "conclusions": "The idea that the mass of dark matter can be converted into a cosmic ray signal has long inspired numerous searches, leveraging its high energy to compete with backgrounds. In XDM, the cosmic ray signal instead arises by converting the {\\em kinetic} energy of dark matter into a cosmic ray signal, but now leveraging a light dark force carrier to yield a large {\\em scattering} cross section, converting dark matter $\\chi$ into its excited state $\\chi^*$. For  $\\delta = m_{\\chi^*}-m_{\\chi}< 2 m_e$, the natural signal is that of an X-ray line, from $\\chi^* \\rightarrow \\chi + \\gamma$. We have shown here that such a signal is observable, and, moreover, is compatible with recent claims of a 3.56 keV line from clusters and M31. In cases with very suppressed dipoles, a primordial excited fraction of dark matter can explain the signal.  Obviously, as a DM explanation of the 3.56 keV line, the scenario is more involved than that of a simple sterile neutrino. Nonetheless, there are good reasons to pursue a second line of thinking. To begin with, the spatial morphology would allow a comparison (with the XDM signal falling more rapidly with radius), and the $\\rho^2$ dependence might alleviate some tensions present with the sterile neutrino model, in particular with Virgo.   The XDM scenario also offers different search strategies - for instance, one could stack a large set X-ray observations of nearby ellipticals selected from the SDSS catalog. Finally, one can look for XDM signals in direct detection experiments, as well as at accelerators in searches for the force carrier $\\phi$. These final strategies are relevant when primordial excited state decay explains the line. Should this signal persist, we then have great hope that additional observations could clarify what is the origin of the 3.56 keV line."
    },
    "1402/1402.0866_arXiv.txt": {
        "abstract": "Exploiting the deep high-resolution imaging of all 5 CANDELS fields, and accurate redshift information provided by 3D-HST, we investigate the relation between structure and stellar populations for a mass-selected sample of 6764 galaxies above $10^{10}\\ M_{\\sun}$, spanning the redshift range $0.5<z<2.5$.  For the first time, we fit 2-dimensional models comprising a single \\sersic\\ fit and two-component (i.e., bulge + disk) decompositions not only to the $H$-band light distributions, but also to the stellar mass maps reconstructed from resolved stellar population modeling.  We confirm that the increased bulge prominence among quiescent galaxies, as reported previously based on rest-optical observations, remains in place when considering the distributions of stellar mass.  Moreover, we observe an increase of the typical \\sersic\\ index and bulge-to-total ratio (with median $B/T$ reaching 40-50\\%) among star-forming galaxies above $10^{11}\\ M_{\\sun}$. Given that quenching for these most massive systems is likely to be imminent, our findings suggest that significant bulge growth precedes a departure from the star-forming main sequence.  We demonstrate that the bulge mass (and ideally knowledge of the bulge and total mass) is a more reliable predictor of the star-forming versus quiescent state of a galaxy than the total stellar mass.  The same trends are predicted by the state-of-the-art semi-analytic model by Somerville et al.  In the latter, bulges and black holes grow hand in hand through merging and/or disk instabilities, and AGN-feedback shuts off star formation. Further observations will be required to pin down star formation quenching mechanisms, but our results imply they must be internal to the galaxies and closely associated with bulge growth. ",
        "introduction": "\\label{obs_sample.sec}  The core dataset used for this work is the deep space-based HST imaging from the CANDELS multi-cycle treasury program \\citep{grogin2011, koekemoer2011}, complemented with redshift information from the 3D-HST grism survey (van Dokkum et al. 2011; Brammer et al. 2012)\\footnote{The AGHAST G141 grism observations within GOODS-North from the GO-11600 program (PI: B. Weiner) are included in the analysis under `3D-HST' (Brammer et al. 2012).}.  In order to determine galaxy-integrated masses and SFRs, we make use of additional multi-wavelength data in the CANDELS/3D-HST fields including space-based photometry from {\\it Spitzer}/IRAC, {\\it Spitzer}/MIPS and {\\it Herschel}/PACS and an array of ground-based facilities (see Skelton et al. 2014 for a detailed description of the $H$-selected multi-wavelength catalogs produced by the 3D-HST team).  The galaxy-integrated properties are derived following identical procedures as \\cite{sw2011b}.  These include stellar masses, based on $U$- to - $8\\mu m$ broad-band SED modeling using population synthesis models from \\cite{bc03} and SFRs derived from a `ladder of SFR indicators'.  The latter method uses detected emission in either UV + PACS for PACS-detected galaxies (Lutz et al. 2011; Magnelli et al. 2013) or UV + MIPS 24$\\mu m$ for MIPS-detected galaxies to compute the sum of the obscured and unobscured SFR. For galaxies lacking an IR detection, the SFR is adopted from the best-fit SED model. We verified that consistent results are obtained when splitting star-forming and quiescent galaxies based on the above measures of star formation activity or when adopting the rest-UVJ diagram (Wuyts et al. 2007; Williams et al. 2009)  \\subsection{HST imaging} \\label{fields_data.sec}   The HST CANDELS observations used for this study comprise high-resolution imaging in five distinct fields: GOODS-South, GOODS-North, COSMOS, UDS, and EGS, covering a total area of 625 arcmin$^2$.  Typical limiting depths in \\H160 are 27.0 mag for CANDELS/wide and 27.7 mag for CANDELS/deep (the central halves of the GOODS fields), respectively.  Additional data used for this work include pre-existing ACS imaging in the GOODS, EGS, and COSMOS fields (Giavalisco et al. 2004; Davis et al. 2007; Koekemoer et al. 2007), plus WFC3 imaging in ERS \\citep{windhorst2011}. The available passbands are $B_{435}$,$V_{606}$,$i_{775}$,$z_{850}$,$J_{125}$,$H_{160}$ for the GOODS fields and $V_{606}$,$I_{814}$,$J_{125}$,$H_{160}$ for the remaining fields. Additionally, $Y_{098}$ imaging was available for ERS as part of GOODS-South and $Y_{105}$ for the regions with CANDELS/Deep coverage.  All imaging used in our analysis was drizzled to a $0\\farcs06$ pixel scale.  For details on the observations and data reduction, we refer the reader to \\cite{koekemoer2011} and \\cite{grogin2011}.  \\subsection{Sample definition} \\label{fields_data.sec}  The main focus of this paper is on the structural shape of stellar mass distributions in a mass-selected sample of galaxies in the redshift range $ 0.5 < z < 2.5$. For this purpose, we only apply a mass cut  $M > 10^{10}M_{\\odot}$ to select our galaxies, well above the mass completeness limit of the $H_{160}$-selected catalogs for all five fields.  Our sample consequently spans a wide range of SFRs, from normal SFGs on the main sequence to quiescent galaxies (QGs) below, that already formed the bulk of their stars.  For selection, we use the SED modeled galaxy-integrated mass estimates, but note that they match well the masses obtained by summing the resolved stellar mass distributions \\citep[see also][]{sw2012}.  In order to determine redshifts for the galaxies in our sample, we use ground-based spectroscopic redshift information whenever available. Otherwise, redshifts are fitted to the combination of 3D-HST grism data and broad-band photometry.  The total sample comprises 6764 galaxies, of which 3839 and 2925 lie within the redshift range $ 0.5 < z < 1.5$ and  $ 1.5 < z < 2.5$, respectively.  In the following, these two redshift ranges are referred to as $z \\sim 1$ and $z \\sim 2$, respectively. ",
        "conclusions": ""
    },
    "1402/1402.5502_arXiv.txt": {
        "abstract": "The  standard cosmological  model, $\\Lambda$CDM, provides an excellent fit  to Cosmic Microwave Background (CMB) data. However, the model has well known problems. For example, the cosmological constant, $\\Lambda$, is fine-tuned to 1 part in $10^{100}$ and the cold dark matter (CDM) particle is not yet detected in the laboratory.   Shanks previously investigated a model which assumed neither exotic particles nor a cosmological constant but instead postulated a low Hubble constant ($H_0$) to allow a baryon density compatible with inflation and zero spatial curvature. However, recent Planck results make it more difficult to reconcile such a model with CMB power spectra. Here we relax the previous assumptions to assess the effects of assuming three active neutrinos of  mass $\\approx5eV$. If we assume a low $H_0\\approx45$kms$^{-1}$Mpc$^{-1}$ then, compared to the previous purely baryonic model, we find a significantly improved fit to the first 3 peaks of the Planck  power spectrum. Nevertheless, the goodness-of-fit is still significantly worse than for $\\Lambda$CDM and would require appeal to unknown systematic effects for the fit ever to be considered acceptable. A further serious  problem is that the amplitude of fluctuations is low ($\\sigma_8\\approx0.2$) making it difficult to form galaxies by the present day.  This might then require seeds, perhaps from a primordial magnetic field, to be invoked for galaxy formation. These and other problems demonstrate the difficulties faced by models other than $\\Lambda$CDM in fitting ever more precise cosmological data. ",
        "introduction": "The $\\Lambda$CDM model is highly successful at fitting the phenomenology of observational cosmology including the CMB and large-scale matter power spectra and these are highly important successes. However, the model suffers from problems at a more fundamental level. First the size of the cosmological constant $\\Lambda$ implies a variety of fine tuning \\citep{Carroll01}. For example, in the early Universe at the end of the inflationary epoch, the ratio of the vacuum energy implied by $\\Lambda$ to the energy in the radiation is 1 part in $\\approx10^{100}$. The vacuum energy of the cosmological constant can be replaced with dark energy whose density can evolve with time and thus alleviate this fine tuning. However, there still remains the coincidence of why the matter and dark energy densities reach equality so close to the present day. The dark matter component of the model also has the problem that the favoured candidate, the neutralino,  is as yet undetected in the laboratory. The lower limits on supersymmetric particle masses such as the s-quark have reached $>1TeV$ at the LHC and almost rule out the Minimal Supersymmetric Model (MSSM) (eg \\citet{buchmueller13}). New limits from direct detection dark matter experiments such as LUX have ruled out a large part of the WIMP mass plus WIMP-nucleon cross-section plane of interest to MSSM \\citep{lux13}. Previous claims of WIMP direct detections  at $\\approx 10GeV$ masses have been comprehensively ruled out by the LUX data. Finally,  the idea  that new supersymmetric particles may exist at masses of a few hundred GeV is difficult to reconcile with the absence of an electron electric dipole moment at present experimental limits \\citep{hudson11, acme13}.  The $\\Lambda$CDM model also has well-known astrophysical problems. In particular, the halo mass function  increases like a power-law towards small scales and looks little like the luminosity or  stellar mass functions of galaxies which exhibit a  characteristic knee in their distribution at around Milky Way size. This has to be addressed by feed-back from supernovae and/or AGN  which keeps the smallest haloes dark \\citep{benson03, baugh05, bower06}. Star-formation feedback is also used to suppress the visibility of sub-haloes in the Milky Way which would otherwise overpredict the number of  satellite galaxies by more than  an order of magnitude. However, these feedback mechanisms may have issues. It has been argued that some of the Milky Way sub-haloes are `too big to fail' and so cannot be simply erased by feedback \\citep{boylan11}. It has also been argued that although a feedback prescription can reproduce the luminosity function of Milky Way satellites their M/L and/or their density concentrations may still be incorrect eg \\citet{zavala09}. Also the known MW and Andromeda satellites may be found in a planar configuration, difficult to reproduce in a merging model like $\\Lambda$CDM \\citep{ibata13}. There have been other claims that dwarf galaxies have cores rather than the cusps predicted by $\\Lambda$CDM \\citep{moore94}. Essentially these are all symptoms of the fact that the top-down structure formation of $\\Lambda$CDM model produces too much power at small scales. However, there are other issues including the lack of merging evident in the evolution of the stellar mass function of even the reddest galaxies.  These problems have led several authors to look at other models such as Warm Dark Matter (WDM)  or Modified Gravity. For example, \\citet{lovell12} have investigated the possibilities of a 1keV sterile neutrino. \\citet{mcgaugh04} also suggested that a model with $\\Omega_\\Lambda=0.97$, $\\Omega_b=0.02$ and $\\Omega_\\nu=0.01$ could fit the early WMAP data. \\citet{angus09} have suggested that a model with three massless active neutrinos and one 11eV sterile neutrino and a cosmological constant gives a good fit to the CMB power spectrum. Here, also motivated by the issues for $\\Lambda$CDM,  we further consider the pros and cons of  Hot Dark Matter (HDM) models. We start with the low-$H_0$ baryonic model of Shanks et al  as an example of the `what you see is what you get' approach in terms of the efforts that have been made to reconcile it to the WMAP and Planck CMB data.  Thus in Section \\ref{sect:baryon_model} we therefore describe the low $H_0$ baryon dominated model of \\citet{shanks85}. In Section \\ref{sect:nu_model} we discuss how a cosmological model that assumes a 5eV mass for each active neutrino species produces a much improved fit to the CMB power spectrum. In Section \\ref{sect:simn} we shall simulate neutrino universes using GADGET2 to assess the usual issues for galaxy formation in neutrino models.  In Section \\ref{sect:discussion} we shall discuss whether primordial magnetic fields might be able to seed galaxy formation in the neutrino model and in Section \\ref{sect:conclusions} we present our conclusions. ",
        "conclusions": "\\label{sect:conclusions}  We have found that a cosmological model with three $\\approx5$eV active neutrinos produce, at least relative to the previous purely baryonic model of \\citet{shanks85}, an  improved fit to the first three peaks of the microwave background power spectrum if a low value of $H_0\\approx45$kms$^{-1}$Mpc$^{-1}$ is assumed.  Even here the model produces a significantly poorer quality of fit than $\\Lambda$CDM. The model further overestimates the amplitude of the fourth, fifth and sixth peaks but this agreement may be improved by smoothing due to lensing or beam profile systematics \\citep{shanks07,whitbourn14a}. Nevertheless, the neutrino model is formally rejected by the CMB data at $\\approx9\\sigma$ significance even when an ad hoc  lensing model is applied. This is significantly worse than the fit achieved by $\\Lambda$CDM. We are here simply recording our view that this neutrino model may be the best that can be done without invoking an exotic new particle or a cosmological constant.   Even ignoring the significantly poorer CMB fit than $\\Lambda$CDM, the main problem with our neutrino model concerns the difficulty in forming galaxies due to the free-streaming of the neutrinos. The neutrino model provides a matter power spectrum with a turnover at $\\approx25$h$^{-1}$Mpc caused by free-streaming which erases fluctuations on small scales. We have concluded that galaxy formation seeds are required for the initial conditions and we have suggested that primordial magnetic fields may provide such seeds. Such fields can produce a steep  matter power spectrum useful for galaxy formation in hot dark matter models. The baryon power spectrum generated by PMF will dominate the power spectrum at small scales. We have noted that such a power spectrum may lead to models more like the monolithic collapse models that the galaxy evolution data favour rather than the merging dominated galaxy formation of $\\Lambda$CDM. We noted that the BICEP2 claim to detect large-scale B-mode polarisation on the CMB may support the existence of PMF. We have also  noted that PMF may already have problems evading Planck non-Gaussianity upper limits (eg \\citet{trivedi14}). PMF at the levels required for galaxy formation in our neutrino model will also be detectable by forthcoming  Planck CMB polarisation  results and other seed mechanisms would have to be sought if  the required primordial magnetic fields are  confirmed to be  ruled out.  There are  other problems with the neutrino model. The value of $H_0$ is low and would require help from a local hole underdensity and other systematic issues with SNIa and Cepheid distances. The matter power spectrum for the model contains baryon acoustic oscillations but these are at too low an amplitude to be compatible with the acoustic peak seen in the  DR11 CMASS galaxy correlation function. If the errors on the correlation function are reliable then this would also present a serious problem for any  neutrino-dominated  model. The 5eV neutrino masses are compatible with SN1987A upper limits on the neutrino mass \\citep{pagliaroli10} but  are already $\\approx5\\sigma$ above the upper limit from tritium  $\\beta$ decay experiments \\citep{aseev11}. Certainly experiments like KATRIN \\citep{katrin01} should be able to detect or reject this mass for the electron neutrino at high significance.    We conclude that while the inclusion of 5eV  active neutrinos can certainly improve the CMB power spectrum fit compared to a baryon dominated model, the model still produces a less good fit than $\\Lambda$CDM and this and the other observational problems we have listed  illustrate the difficulty in finding acceptable alternatives to the standard $\\Lambda$CDM cosmology."
    },
    "1402/1402.3504_arXiv.txt": {
        "abstract": "A simple method has been devised to reduce the computing time needed to calculate the electric field of radio pulses emitted from air showers. It relies on an approximate description of the patterns of the radio pulses produced by two different mechanisms: the excess charge and the geomagnetic deviation of the charged particles. The amplitudes of the electric field due to each of these mechanisms must be calculated in Monte Carlo simulations but only at a few positions along a line at ground level. The amplitudes at other positions are deduced using the projected Cherenkov cone on the ground. The polarization patterns of the electric field due to each mechanism are different and are taken from relatively well known properties of the radio emission in air showers. The results of this approach, that can be obtained with modest computing resources, are shown to agree at a few percent level with full Monte Carlo simulations. ",
        "introduction": "In the last few years there has been a flurry of activity to explore the potential of the radio technique to study ultra high energy cosmic rays (UHECRs) \\cite{Huege_ICRC13}. The technique is attractive because it requires detectors which are relatively cheap and is not limited by a small duty cycle, unlike the fluorescence technique \\cite{Nagano-Watson}. Several experimental initiatives such as AERA \\cite{AERA} and EASIER \\cite{EASIER} in coincidence with the Pierre Auger Observatory \\cite{Auger}, CODALEMA \\cite{CODALEMA}, LOPES \\cite{LOPES}, TREND \\cite{TREND}, LOFAR \\cite{LOFAR} and Tunka-Rex \\cite{Tunka-Rex_ICRC13},  have been exploring emission from air showers mainly in the 30-80 MHz frequency range, studying the relation of the pulses to composition \\cite{LOPES_muons_radio,LOFAR_Xmax} and trying to understand the lateral distribution of the emission \\cite{Allan_Nature,LOPES_LDF,AERA,LOFAR}. In addition the fortuitous detection of pulses from air showers with the ANITA balloon flown detector \\cite{ANITA_UHECR} has revealed that the pulses from these showers extend to the GHz regime, with more experimental evidence also obtained by the CROME experiment \\cite{CROME}. These studies are considered of primordial importance to explore the viability of the radio technique either to provide an alternative design of future air shower arrays covering areas in the range of 10,000~km$^2$ and above, or to use them as complementary detectors that could help to constrain the composition of the cosmic rays on a shower-by-shower basis.  On the phenomenological side there has also been a lot of progress. Calculations have been made based on the macroscopic treatment of the current densities that arise in these showers \\cite{Scholten_MGMR,Scholten_EVA}, and also on more detailed simulations that use the superposition principle to calculate the emission adding the contributions from single particle sub-tracks \\cite{REAS3,SELFAS,ZHAireS,COREAS}. The results of these calculations, that reproduce pulses extending into the GHz regime as detected by ANITA \\cite{ZHAireS_ANITA}, are now in reasonably good agreement \\cite{Huege_ARENA12} and they have been successfully tested against data \\cite{LOPES_COREAS,ZHAireS_ANITA,AERA,AERA_pol,Huege_ICRC13}. As a result there is much more confidence in the calculations and it is believed that it is now possible to infer properties of the shower such as direction, energy and shower maximum using only measurements of the radio pulses in antennas scattered at ground level \\cite{Huege_radio_mass,Konstantinov_JCAP,deVries_radio_mass,Scholten_Xmax,Huege_ICRC13}.  The study of the capabilities of these detectors as well as the reconstruction of cosmic ray shower parameters in detail requires full simulations of the shower and the radio pulses at a large variety of locations what becomes quite demanding of computing resources. This has become a natural obstacle for progress since these simulations have to be performed for a wide range of energies, shower geometries and primary particle compositions what inevitably is a very lengthy and arduous task. Needless to say that a comprehensive parameterization of the patterns of the electric field at ground level would be most welcome as an alternative. The approach presented here can be considered a step forward in this direction, since it can also simplify the creation of parameterizations.  In this work we have devised a phenomenological technique to reproduce the pattern of the radio pulses at ground level under most circumstances. The technique relies on standard approximate assumptions \\cite{Scholten_MGMR} and experimental observations of the characteristics of the radiation \\cite{Allan,Prescott,CODALEMA,AERA_pol}. These are mainly due to two different radiation mechanisms, namely the excess negative charge predicted by Askaryan in the 1960s, the Askaryan mechanism \\cite{Askaryan}, and that induced by the interactions of the particles with the magnetic field of the Earth that induces a net current perpendicular to the shower axis, the geomagnetic mechanism \\cite{Kahn-Lerche}. We present here a method that exploits these general ideas and subsequently test the level of accuracy that it provides comparing the obtained results directly with the detailed simulations. The devised method reduces the calculation of the pattern at ground level to the pattern along a single line which can be sampled at a few points. This has a drastic effect on the required computing resources and has many applications for both the reconstruction of events at a given experimental facility or for design studies of future experimental arrangements.  This paper is organized as follows. In Section \\ref{sec:basics} we briefly review the main properties of radio emission in atmospheric cosmic ray showers. In Section \\ref{sec:model} we describe a simple model to obtain the electric field induced by ultra-high energy cosmic ray showers in the atmosphere. The model only needs as input the separate amplitudes of the electric field induced by the geomagnetic and Askaryan contributions at a few positions on the ground, as obtained in Monte Carlo simulations. It then assumes a radial symmetry of the geomagnetic and Askaryan amplitudes and exploits the expected polarization of each mechanism to obtain the net electric field at any position on the ground. In Section \\ref{sec:results} we compare the electric field calculated with this approach with that obtained in full Monte Carlo simulations. Finally, in Section \\ref{sec:conclusions} we conclude the paper and discuss future prospects. ",
        "conclusions": "\\label{sec:conclusions}   We have shown that using standard assumptions on the polarization and symmetry properties of the electric field induced by the geomagnetic and Askaryan mechanisms, we are capable of estimating with great precision the electric field at any position on the ground by using as input simulations of the electric field at just a few positions. We have shown that in the majority of cases the estimates are within just a few percent of those obtained by performing time consuming full Monte Carlo simulations for all antennas. More than a simple model, this can be seen as an approach that takes into account all the complexities of the shower radio emission in a simple manner, namely by treating the two main emission mechanisms separately and looking for symmetries that do not apply to the total net electric field. Although not shown in this work, we expect our method to work at other energies and for other primaries besides the ones used here. We have also discussed the sources of inaccuracies that increase the differences between the predictions and full simulations. By refining the assumptions described in Section \\ref{sec:assumptions} to include the discrepancies found, we could not only further improve the accuracy of the model, but also better understand the main characteristics of the radio emission in the process.  The most clear advantage of using the method described in this work to predict the electric field is a significant reduction in computing time. As an illustration of what is possible we show in Fig.~\\ref{fig:map} a two-dimensional map of the electric field on the ground for which we used 36 antenna positions East of the core as input\\footnote{This number of antennas has not been yet optimized and can be significantly further reduced.}. The total computing time used for the input simulations was around 10 h in a single processor. The input was used to estimate the field at antennas at every 2 m in a rectangular grid, totaling around 200 thousand antennas. This process takes just a few minutes. Our approach could also prove to be a faster alternative to radio signal interpolation techniques, which need a much larger number of simulated antenna positions on the ground to obtain similar results.  This approach may also prove to be a useful tool to simplify the creation of electric field parameterizations. As already discussed, the actual input are just two curves of $|\\vec{E}|$ as a function of distance $R$ to the core along a reference line, one for the Askaryan contribution and another for the geomagnetic one. So a parameterization of these two curves along a single arbitrary reference line as a function of the shower parameters (in a given experiment site), namely the energy $E_0$, shower maximum $X_{\\rm max}$ and arrival direction could be used to create precise estimates of the electric field at any point on the ground. By simulating a large number of showers to create such parameterizations, one could obtain distributions for the parameters, making possible the inclusion of shower to shower fluctuations.  We expect the proposed method to be very useful in many applications, including the study of the characteristics of the radio lateral distribution function, such as its asymmetries and their dependence on shower parameters, and experimental reconstructions that are based on Monte Carlo simulations of the radio emission, such as the $X_{\\rm max}$ reconstruction of the LOFAR experiment \\cite{LOFAR_Xmax}. It could also be used as a test bench for new, more robust theoretical predictions for the characteristics of each emission mechanism."
    },
    "1402/1402.4324_arXiv.txt": {
        "abstract": "{The B-mode polarization of the cosmic microwave background (CMB) radiation is an excellent information channel for the detection of relic gravitational waves. However, the detection is contaminated by the B-mode polarization generated by some other effects. In this paper, we discuss the contaminations caused by the cosmological birefringence, which converts the CMB E-mode to the B-mode, and forms the effective noise for the detection of gravitational waves. We find that this contamination is significant, if the rotation angle is large. However, this kind of B-mode can be properly de-rotated, and the effective noises can be greatly reduced. We find that, comparing with the contaminations caused by cosmic weak lensing, the residual polarization generated by the cosmological birefringence is negligible for the detection of relic gravitational waves in the CMB. } ",
        "introduction": "The temperature and polarization anisotropies of the cosmic microwave background (CMB) radiation contain fruitful cosmological information, and play crucial roles in the determinations of various parameters in modern cosmology. In particular, the CMB anisotropies provide the unique observational way to detect the relic (primordial) gravitational waves, which was inevitably produced in the very early Universe \\cite{grishchuk1979,starobinsky1980}. Relic gravitational waves left the observable imprints in all the CMB power spectra, including the TT, TE, EE and BB information channels. Limited by the cosmic variance, if the tensor-to-scalar ratio $r<0.05$ \\cite{turner1994,zhao2009}, the TT, TE and EE channels become useless, and we can only detect it in the B-mode polarization \\cite{zaldarriaga1997,kamionkowski1997}, which is one of the most important goals for the CMB efforts \\cite{task,krauss2010}. {\\tb {The current satellite observations, including those of the WMAP \\cite{wmap9} and Planck missions \\cite{planck}, are yet to detect a definite signal of relic gravitational waves. However, the recent observations of the ground-based experiments, such as the BICEP1, SPTPOL, POLARBEAR, ACTPOL telescopes, have given some interesting results for the B-mode polarizations \\cite{bicep1,sptpol,polarbear,actpol}. In particular, the BICEP2 team released their recent data, and claimed the discovery of relic gravitational waves with the tensor-to-scalar ratio $r=0.20^{+0.07}_{-0.05}$, and $r=0$ disfavored at $7.0\\sigma$ \\cite{bicep2}. However, it has been shown that this observation is also consistent with the polarized radiation emitted by the poorly-understood interstellar dust \\cite{debate1,debate2,liuhao}. In order to distinguish them, one should measure the B-mode polarization at the different frequency channels by BICEP2 or some other experiments (such as the Planck satellite, BICEP3, EBEX, QUBIC and so on). Nevertheless, these results encourage us to put the gravitational-wave detection through the CMB polarization as a highest priority for the next generations of the CMB experiments \\cite{task,krauss2010}. In addition, we should mention that it is important to cross-check the CMB results by the other large-scale observations. For example, in the previous works \\cite{c1,c2}, the authors found that the relic gravitational waves can distort the shapes of galaxies through the effects of tidal fields of the large-scale structure and the gravitational lensing, and form the B-modes in galaxy shape correlations, which can also be used to detect the signal of gravitational waves \\cite{d1,c3}. In \\cite{c3}, the authors discussed in detail the possibility of the gravitational-wave detection by using the cross-correlation between shear and CMB B-mode polarization, and found that it is quite possible to confirm or falsify the BICEP2 results, if the next-generation surveys beyond EUCLID, WFIRST, and LSST are considered. }}     In the standard cosmological model, up to the first-order perturbations, {\\tc {both the vector and tensor perturbations \\cite{zaldarriaga1997,kamionkowski1997} have the possibilities to generate CMB B-mode. However the vector perturbations usually decay quickly in the expanding universe and can be ignored in CMB physics. Hence significant primordial CMB B-mode can only be generated by the tensor perturbations, i.e., the relic gravitational waves. In principle, such signals for the relic gravitational waves}} cannot be polluted by the primordial density perturbations, which is the reason why the method of detecting gravitational waves through CMB B-mode is clean. However, in the real Universe, the detection ability of the method is limited by various contaminations. In addition to the foreground radiations, various systematic errors, the E-B mixture caused by the incomplete sky survey, and the instrumental noises in the real observations, some other effects can also naturally generated the B-mode polarization. The well-studied one is the cosmic weak lensing, which mixed the E-mode and B-mode polarizations, and forms a nearly multipole-invariant BB power spectrum \\cite{zaldarriaga1998,lewis2006}.  Another effect to produce the B-mode polarization is the so-called cosmological birefringence, which can be caused by the possible coupling between the electromagnetic field and the scalar field (which may or may not be identified as the dark energy) through the Chern-Simons term \\cite{carroll,kamionkowski1998,other-cpt,li2009}. The cosmological birefringence generates a frequency-independent rotation of the linear polarization of the CMB photons when they propagate over the cosmological distances. So the B-mode polarization can be naturally converted from the E-mode, even if it is absent in the early universe. This phenomenon resulted from the Lorentz and $CPT$ violations in the electrodynamics provides an effective method to test fundamental symmetries of nature and attracted many interests. {\\tc {It is similar to but different from the Faraday rotation caused by the cosmological magnetic field \\cite{Faraday1,Faraday2,Faraday3}, where the rotation of the polarization depends on the frequency of the photon. In this paper the ``cosmological birefringence\" only means the frequency-independent rotation, as in Refs. \\cite{carroll,kamionkowski1998}. For this phenomenon,}} the rotation of the CMB photons is only quantified by the rotation angle $\\alpha$. Numerous works have constrained it by using the current CMB data, and show some evidences of the nonzero result \\cite{feng2006,other-constraints,xia2012}. The current tightest constraint comes from the data analysis in \\cite{xia2012}, where the author found that ${\\alpha}=-2.28\\pm 1.02  ~{\\rm deg}$ (1$\\sigma$) when considering the seven-year WMAP, BOOMERanG 2003 and BICEP data. This follows that $|\\alpha|<4.32 ~{\\rm degree}$ in the $2\\sigma$ confidence level. However, we should remember that this result is still in debate. For instance, in the same paper, the author also found that $-1.34 < \\alpha < 0.82~{\\rm degree}$ at $95\\%$ confidence level, if adding the QUaD polarization data.  In this paper, we shall investigate the CMB B-mode polarization produced by the cosmological birefringence, and focus on its influence on the detection of relic gravitational waves. We find that this B-mode could be quite large, if the rotation angle is close to the current upper limit value. If considering this B-mode as a new effective noise, the gravitational-wave detection in the CMB is impossible when $r<0.0014$. So, it is important to remove this contamination for the future observations. In our discussion, we propose a method to de-rotate it by utilizing the statistical properties of the E-mode polarization, and those of the estimator of $\\alpha$ parameter. We find that, if considering the de-rotating, the residual B-mode polarization becomes very small. Comparing with contamination of the cosmic weak lensing, the residuals become negligible for the detection of gravitational waves.  The outline of this paper is as follows. In Sec. \\ref{sec2}, we briefly introduce the cosmological birefringence and the rotated CMB linear polarizations. In Sec. \\ref{sec3}, we discuss the method to de-rotate the CMB B-mode polarization, and the influence on the gravitation-wave detection. In Sec. \\ref{sec4}, we summarize the main results of this paper. ",
        "conclusions": "}  The cosmological birefringence caused by the Chern-Simons coupling of the cosmic scalar field to the electromagnetic field has the possibility to rotate the polarization planes of the CMB photons when they propagate from the last scattering surface to us and to convert a part of  E-mode polarization to the B-mode polarization. Such kind of rotated BB power spectrum at late time forms a new contamination for the detection of relic gravitational waves in the CMB. If the rotation angle $\\alpha$ is close to the current upper limit value, we find that the gravitational-wave detection is limited by this new noise if the tensor-to-scalar ratio is smaller than $0.0014$.  In this paper, we suggest the method to partly reconstruct and subtract the rotated B-mode polarization by utilizing the statistical properties of the estimators of E-mode coefficients $E_{\\ell m}$ and the rotation angle $\\alpha$. We find that, if this de-rotating is done by considering the noise level of the CMBPol mission or the better experiments, the residual BB power spectrum can be reduced by more than two orders, even if the largest rotation angle is considered. The residuals are much smaller than the de-lensed BB power spectrum caused by the cosmic weak lensing, and become negligible for the detection of gravitational waves.    ~  ~  ~  {\\it Acknowledgments:} W.Z. is supported by project 973 under Grant No.2012CB821804, by NSFC No.11173021, 11322324 and project of KIP of CAS. M.L. is supported by Program for New Century Excellent Talents in University and by NSFC under Grants No. 11075074.   \\appendix"
    },
    "1402/1402.4865_arXiv.txt": {
        "abstract": "We use $N$-body simulations as well as analytical techniques to study the long term  dynamical evolution of stellar black holes (BHs) at the Galactic center~(GC) and to put constraints on their number and  mass distribution. Starting from  models that have not yet achieved a state of collisional equilibrium, we find that time scales associated with cusp regrowth can be longer than the Hubble time. Our results cast doubts on standard models that postulate high densities of BHs near the GC and motivate studies that start from initial conditions which correspond to well-defined physical models. For the first time, we consider the distribution of BHs in a dissipationless  model for the  formation of the Milky Way nuclear cluster~(NC), in which massive stellar clusters merge to form a compact nucleus. We simulate the consecutive merger of $\\sim 10$ clusters containing an inner dense sub-cluster of BHs. After the formed NC is evolved for $\\sim 5$~Gyr, the BHs do form a steep central cusp, while the stellar distribution maintains properties that resemble those  of the GC NC. Finally, we  investigate the effect of BH perturbations on the motion of the GC S-stars, as a means of constraining  the number  of the perturbers. We find that reproducing the quasi-thermal character of the S-star orbital eccentricities requires $\\gap 1000$ BHs within $0.1~$pc of Sgr~A*. A dissipationless formation scenario for the GC NC is consistent with this lower limit and therefore could reconcile the need for high central densities of BHs (to explain the S-stars orbits), with the ``missing-cusp'' problem of the GC giant star population. ",
        "introduction": "Massive nuclear clusters~(NCs) are observed at the center of many galaxies, over the whole Hubble sequence. The frequency of nucleation among galaxies less luminous than $\\sim10^{10.5}~L_\\odot$ is close to $90~\\%$  as determined by ACS HST  observations of galaxies in the Virgo and  Fornax galaxy clusters~\\citep{carollo,BSM:02,Cote,turner+12}. The study of NCs is of great interest for our understanding of galaxy formation and evolution as indicated by the fact that a number of fairly tight correlations are observed between their masses  and  global properties of their host galaxies such as  velocity dispersion and bulge mass~\\citep{F:06,WH:06,GS:09,SG12,Nathan}. Intriguingly, similar scaling  relations are obeyed by  massive black holes~(MBHs) which are predominantly found in massive galaxies that, however, show little evidence of nucleation~\\citep[e.g.,][]{GS:09,NW12}. The existence of such correlations  might indicate a direct link among large galactic spacial  scales and the much smaller scale of the nuclear environment, and suggests that NCs contain information about the processes that have shaped the central regions of their host galaxies.  How NC  formation takes place at the center of galaxies is still largely debated~\\citep[e.g.,][]{Hartmann,gnedin+13,carlberg+hartwick14,mpl14}. Relatively recent work has shown that ``dissipationless''  models can reproduce without obvious difficulties the observed properties~\\citep{turner+12} and scaling relations \\citep{me13} of NCs. In these models a NC forms through the inspiral of massive stellar clusters into the center due to dynamical friction where they merge to form a compact nucleus~\\citep[e.g.,][]{TOS,CDM08,capuzzo93}. Alternatively, NCs could have formed locally as a result of radial gas inflow into the galactic center accompanied  by efficient dissipative processes~\\citep{Sch08,Milos04}. Naturally,  dissipative and dissipationless processes are not exclusive and both could be important for the formation and evolution of NCs~\\citep{Hartmann,AM12-2,de-lorenzi+13}.   The Milky Way NC, being only $8~$kpc away, is currently the only NC that can be resolved in individual stars and for which a  kinematical structure and density profile can be reliably determined~\\citep{gen+10}. This offers the unique possibility to resolve the stellar population, to study the composition and dynamics close to a MBH and put constrains on different NC formation scenarios. The Milky Way NC has an estimated mass of $\\sim 10^{7}M_{\\odot}$ \\citep{LZM,S09}, and it hosts a massive black hole  of $\\sim 4\\times 10^6 M_\\odot$~\\citep{Genzel03,Ghez08,Gil} whose gravitational potential dominates over the stellar cusp potential out to a radius of roughly $3$pc - the MBH radius of influence. A handful of other galaxies are also known to contain both a NC and a MBH, which typically have comparable masses~\\citep{seth}. Population synthesis models suggest that roughly $80\\%$ of the stellar mass in the inner parsec of the Milky Way is in ($>5~$Gyr) old stars~\\citep{Ol11} although the light is dominated by the young stars. This appears to be typically the case also in most NCs observed in external galaxies \\citep{Rossa}.  Over the last decades observations of the Galactic NC have  led to a number of puzzling discoveries. These include: the presence of a young population of stars (the S-stars) near Sgr~A* in  an environment extremely hostile to star formation~\\citep[paradox of youth,][]{Morriss1993,sch02}; and a significant paucity of red giant stars in the inner half a  parsec~\\citep[conundrum of old age,][]{M:10}. Number counts of the giant stars at the Galactic center~(GC) show that their visible distribution is in fact quite inconsistent with the distribution of stars expected for a dynamically relaxed population  near a dominating Keplerian potential~\\citep{BSE,D:09,B:10}: instead of a steeply rising \\citet{BW:76} cusp, there is a $\\sim 0.5$ pc core. The lack of a Bahcall-Wolf cusp in the giant distribution  casts doubts on  dynamical relaxed, quasi-steady-state models  of the GC which postulate a  high central density of stars and stellar black holes~(BHs). In these models the central distribution  of stars and BHs is determined by just a handful of parameters: the MBH mass; the total density outside the relaxed region; the slope of the initial mass function~\\citep[IMF,][]{Merritt-book}. Given the unrelaxed form of the density profile of stars, making predictions about  the distribution of the stellar remnants becomes a much more challenging, time-dependent, problem susceptible  to the initial conditions and to the (yet largely unconstrained) formation process of the NC~\\citep{AM12}.    Understanding the distribution of the ``stellar remnants'' in systems similar to the Milky Way's NC  is crucial in many respects. Examples include randomization of the S-star orbits via gravitational encounters~\\citep{Perets2009}, warping of the young stellar disk~\\citep{kocsis+tremaine11}, and formation of X-ray binaries~\\citep{muno+05}. Stellar nuclei similar to that of the Milky Way are also the location of astrophysical processes that are potential  gravitational wave~(GW) sources both for   ground and space  based laser interferometers. These include the merger of compact object binaries  near MBHs~\\citep{me+perets12}, and the capture of BHs by MBHs,  called ``extreme mass-ratio inspirals\"~\\citep[EMRIs][]{AS12}. The efficiency of these dynamical processes and  rate estimates for GW sources are very sensitive to the number of BHs near  the center. Therefore, a fundamental question  is whether given a prediction for the initial distribution of stars and BHs, the system is old enough  that the heavy remnants  had time to relax and segregate to the  center of the Galaxy.     Motivated by the above arguments,   we consider the long-term evolution of BH populations at the center of galaxies, starting from different assumptions regarding their initial distribution. Since the stellar BHs at the GC are not directly detected, time-dependent numerical  calculations, like the ones presented below, are crucial for understanding and making predictions about the distribution  of stellar remnants at the center of galaxies.  In Section~\\ref{sec-mass-seg} we explore the evolution of  models in which stars  and BHs follow initially  the same spatial distribution which is far from being in collisional equilibrium. Contrary to some previous claims~\\citep{preto-amaro10}, we find that in these models the time to regrow a cusp in both the BH and the star distribution  is longer than the age of the Galaxy. For realistic number fractions of BHs, our simulations demonstrate that over the age of the Galaxy the presence of a heavy component has little effect on the evolution of the stellar component.   In Sections~\\ref{timescales},~\\ref{ics} and ~\\ref{results} we discuss the evolution of BHs in a globular cluster merger model for NCs. We present the results of  direct $N$-body simulations of the merger of  globular clusters containing two mass populations: stars and BHs. These systems were   in an initial   state of mass segregation with the BH population concentrated toward the cluster core. Each cluster was  placed on a circular orbit with galactocentric radius of $20~$pc in a $N$-body system containing a central MBH. We find that the inspiral of massive globular clusters in the center of the Galaxy constitutes an efficient source term of BHs in these regions. After about ten inspiral events the BHs are highly segregated to the center. After a small fraction of the nucleus relaxation time (as defined by the main stellar population) the BHs attain a nearly-steady state distribution; at the same time the stellar density profile exhibits a $\\sim 0.2$pc core, similar to the  size of  the core in the distribution of  stars at the GC. Our results indicate that standard models, which assume the same initial phase space distribution for BHs and stars, can lead to misleading  results regarding the current dynamical state of the Galactic center.  We discuss the implications of our results in Section~\\ref{discussion}. In particular, we show that in order to reproduce the quasi-thermal form of the observed eccentricity distribution of the S-star orbits, about $1000~$BHs should be present inside $\\sim 0.1$pc of Sgr~A*. This number appears to be consistent with the number of BHs  expected in a model in which the Milky Way NC formed trough the orbital decay and merger of about $10$ massive clusters.  Our main results are summarized in Section $7$. ",
        "conclusions": "\\label{discussion}   \\subsection{Comparing to  the properties of the Galactic Nuclear Cluster}  Star counts using adaptive optics spectroscopy and medium-band imaging have shown that the red giants at the Galctic center, the only old stars that can be resolved in these regions, have a flat projected surface density profile close to Sgr~A*~\\citep{BSE,D:09}. The core in the red giant  population extends out to approximately $0.3-0.5~$pc from the center. However, due to the effect of projection, it is difficult to constrain  core size and  three-dimensional spatial density profile, which could be slowly rising but even declining toward the center.  It is possible that a  cusp in the lower mass stars is present and that the observed core is the result of a luminosity function that changes within $0.5~$pc. This could be due to physical collisions before or during the giant phase~\\citep{bailey+davies99,dale+09}; or tidal interactions between  stars and the  central MBH~\\citep{davies+king05} -- mass removal can make the luminosity that a star would otherwise reach at the tip of the red-giant phase considerably fainter and prevent  these objects from evolving to become observable. While these models are possible  they seem not to fully explain the observations~\\citep{dale+09}. Thus, it is important to consider the possibility that the  red giants are indeed representative of the unresolved low-mass main-sequence stars, and that the the density core is a consequence  of  the NC formation history combined to a long nuclear relaxation timescale~(Section~\\ref{sec-mass-seg}). On this basis we can directly compare the predictions of theoretical models to the kinematics and mass distribution inferred from observations of the giant stars at the GC and draw conclusions about possible formation  mechanisms for the Milky Way NC.  \\subsubsection{Mass distribution} Figure~\\ref{DEN1} shows that the merger of about $10~$clusters results into a compact nucleus with a stellar density profile that declines as $n \\sim r^{-1.5}$ outside $\\sim0.5~$pc and  flattens to $n \\sim r^{-1}$ within this radius. More precisely, we consider two definitions of the core radius: (1) the projected radius, $r_c$, at which the surface density falls to one-half its central value; (2) the break radius, $r_b$, at which the density profile transits from the inner law~($n \\sim r^{-1}$) to the outer density law~($n \\sim r^{-1.5}$). We find $r_c = 0.51$pc and $r_b=0.48$pc for Model~B, and $r_c = 0.61$pc and $r_b=0.59$pc for Model~C. Collisional relaxation occurring during the post merger simulation of Model~B reduces the size of the core with time, while inner and outer density slopes remain roughly unchanged. At the end of the simulation, after $\\sim 0.3T_{\\rm r_{\\rm infl}}$,  we find $r_c = 0.32$pc and $r_b= 0.18$pc. The size of the region where the density  transits to a shallow profile in our models appears to be comparable to the  extent of the core inferred from number counts of the red giant stars at the GC. { The exact extent of the core region (and how far our models are from their steady state) is determined by a number of factors that depends on the assumptions made in the $N-$body initial conditions. For example, the size of the density core could be made larger or smaller depending on the initial degree of cluster evolution and the number fraction of BHs. However,  we note that the presence of a core in the final density distribution seems to be a quite robust outcome of a merger  model for NCs. For example, the single mass component simulations of \\citet{AM12-2} also produced a final NC model with a core of size $\\sim 1$pc, somewhat  similar to what we find in this paper. The absence of a Bahcall-Wolf cusp is naturally explained in these models, without the need for fine-tuning or unrealistic initial conditions.  Our simulations result in a final density profile having nearly the same power-law index beyond~$0.5$pc as observed~\\citep[$\\Sigma(R)~R^{-1}$,][]{haller+96}. The slope index in the inner $\\sim 0.5~$pc of our model, $\\gamma \\sim 1$, is also consistent with what obtained from the surface brightness distribution of stellar light within the inner $1''$ of Sgr A*~\\citep{yz+12}, but appears only  marginally consistent  with what inferred from number counts of the red giant stars. Slope indexes in the range ~$ -3 \\lesssim \\gamma \\lesssim 0.8$ are consistent with what derived from observations of the giants, although negative or nearly-zero values, corresponding to centrally-decreasing  or flat densities respectively, are preferred \\citep{M:10,do+13}.   \\subsubsection{Kinematics} The radial profile of the velocity anisotropy of the NC could potentially provide useful constraint on its formation. Kinematic modeling of proper motion data derived from the dominant old population of giants, reveals a nearly spherical central cluster exhibiting slow, approximately solid-body rotation, of amplitude $1.4~{\\rm km s^{-1}arcsec^{-1}}$~\\citep{trippe,S09}. Kinematically, the central cluster appears isotropic, with a mildly radial  anisotropy at $r\\lesssim 0.1~$pc and slightly tangentially anisotropic for $0.1~{\\rm pc}\\lesssim r\\lesssim~1~{\\rm pc}$; In the radial range $1''-10''$, the late-type stars are observed to have a mean projected anisotropy of $\\langle 1-{\\sigma_T^2/\\sigma_R^2\\rangle=-0.12^{+0.098}_{-1.05}}$~\\citep{S09}. \\citet{do+13} found $\\beta=0.01^{+0.35}_{-0.34}$ within $0.5~$pc.  Our models are characterized a generally flat anisotropy profile with $\\beta \\approx -0.5$ at the end of the inspiral simulations and $\\beta \\approx 0$ at the end of the post-merger evolution. Although such values are consistent with observations, we believe that future and better kinematic data that extend outside the inner parsec will be necessary in order to provide better constraints on this scenario.  We note that due to our assumption of no preferential direction of inspiral, the merger remnants in our simulations showed no significant net rotation. Recent observations of the NC in our Galaxy suggest instead a significant rotation on parsec scale \\citep{S14}; this might be reconciled with a cluster merger origin for the Galactic NC if, for instance, the clusters were originally dragged down into the Galactic disk plane (where they experience an a greater dynamical friction force) and transported into the central region of the Galaxy where they then accumulated to form a dense nucleus which will then appear to rotate in the same sense of the Galaxy.    \\subsubsection{Kinematically cold sub-structures} \\citet{Feldme+14}  found indications for a substructure in the Galactic NC that is rotating approximately perpendicular to the Galactic rotation with $\\sim30~{\\rm km~s^{-1}}$ at a distance of $\\sim20\"$ or $0.8~$pc from Sgr A*. In addition they found an offset of the rotation axis from the photometric minor axis and argue that this hints to infalling clusters.  We look for kinematic substructures in our models by using  the Rayleigh \\citep[dipole,][]{rl-19} statistics $R'$ defined as the length of the resultant of the unit vectors $l_i,~i=1,...,N,$ where $l_i$ is perpendicular to the orbital plane of the $i^{th}$ particle and $N$ is the number of particles  (a total of 5715 per cluster). For each merged cluster we computed $R'$ over the entire course of the inspiral simulation. For a fully isotropic distribution we expect $R' \\sim \\sqrt{N}$, while $R' \\sim N$ if the orbits are correlated.  As an example, Figure~\\ref{ksub} gives  $R(=R'/N)$ as a function of time (scaled to the relaxation time of a Milky Way like nucleus).  Initially, after a cluster reaches the center the orbits are strongly correlated and, as expected, $R\\sim1$, i.e., the stars from the infalling clusters  distribute into a thin disk configuration initially. Due to two-body relaxation, and due to the perturbing effect of the later infalling clusters, $R$ decreases  with time and approaches a value more consistent with isotropy. Figure~\\ref{ksub} shows that all clusters maintain some degree of anisotropy during the entire corse of the simulation. But the orbits of  the stars from the first 7 clusters are almost completely isotropic by the end of the simulation. The orbits of stars from the last four/five infalling clusters are still largely correlated after $\\sim 3~$Gyr of evolution. A linear fit to the data ($R$~vs~$t$) for  the last decayed cluster gives: $R(t) = -0.11\\times t/{\\rm Gyr}+0.96$, so that it would take about $\\sim 10$~Gyr for the stars to achieve a nearly isotropic distribution. These results are consistent with those of \\citet{mp13}, who found that it takes a time of order the relaxation time of the nucleus to fully randomize an initially cold disk.   \\begin{figure}\\centering \\includegraphics[width=0.355\\textwidth,angle=270]{Figure18.ps} \\caption{ Evolution of the Rayleigh parameter $R=R'/N$ that measures the degree of randomness of the orbital orientation of  stars transported to the center by each of the 12 infalling clusters of model B.  The hatched region shows the $90\\%$ confidence bands expected for a fully random distribution of orbital orientations of 5715 stars. At the end of the 12th inspiral event significantly cold kinematic substructures are present in the NC model. } \\label{ksub} \\end{figure}    \\begin{figure}\\centering \\includegraphics[width=0.355\\textwidth,angle=270]{Figure19.ps} \\caption{Cumulative number of BHs predicted for a Milky Way nucleus as function of radius and at the different time intervals: $t=(0.09,~\\textcolor{red}{0.18},~\\textcolor{green}{0.27},~\\textcolor{blue}{0.36})\\times T_{r_{\\rm infl}}$, for the post-merger simulation discussed in Section~\\ref{relax}. Dot-dashed curves were obtained by carrying out regression fits of $\\log N_{\\rm bh}$ to $\\log r$ in the radial range $(r_1 , r_2 )$, with $N_{\\rm bh}(< r_1) = 5$  and $N_{\\rm bh}(< r_2) = 25$ (see text for details). } \\label{msclu2} \\end{figure}    \\subsection{Extreme-mass-ratio inspirals} The inspiral of compact remnants into a MBH represents one of the most promising sources of gravitational wave radiation detectable  by space-based laser interferometers~\\citep{AS12}. Event rates for such extreme mass-ratio  inspirals are generally estimated under the assumption that the BHs had enough time to segregate and form a steep central cusp, $n \\sim r^{-2}$, near the MBH. Such dynamical models  predict inspiral rates per galaxies of $\\sim 250~{\\rm Gyr^{-1}}$~\\citep{HA06}. Models that include an initial parsec scale core can result in much lower central BH densities than in the steady state models, and imply rates as low as $1-10~{\\rm Gyr^{-1}}$~\\citep{M:10}. EMRI event rates could be also severely suppressed by the Schwarzschild barrier which limits the ability of stars to diffuse to high eccentricities onto inspiral orbits~\\citep{MAMW2011}.  It must be stressed that it is difficult to draw any conclusion about EMRI event rates from the models discussed in the literature  because of the significant uncertainties in the underlining assumptions. Our simulations cast  further doubts on results obtained from idealized time-dependent models which relay on the assumption that  BHs and stars  have initially the same spatial distribution. For example, based on the models discussed in Section~\\ref{sec-mass-seg} and in \\citet{M:10} and \\citet{AM12}, the presence of a core in the old population of stars at the GC would imply very low central densities of BHs and EMRI event rates. However, the initial conditions adopted in the simulation of Section~\\ref{sec-mass-seg} were quite artificial and not motivated by any specific physical model. We have shown that in a merger model for the formation of NCs, the resulting distribution of stellar remnants partially reflects their distribution in their parent clusters just before they reach the center of the galaxy. Thus, different, possibly more realistic, initial conditions would produce rather different central BH densities;  in these models EMRI rates could be as large as (or higher than)  those obtained in the steady state models even in the presence of a core in the stellar distribution.    In order to determine the number of BHs in a Milky Way like nucleus predicted by the inspiral simulations we used the scaling relation given by Equation~(\\ref{bh-sc}). We are interested in the number of compact remnants inside $\\approx 0.01$~pc, as these are the only BHs that can generate EMRIs. Since  at these radii the number of BH particles  in the $N-$body model  is small, we carried out regression fits of $\\log N_{\\rm bh}$ to $\\log r$ at larger radii, and extrapolate inward in order to get $N_{\\rm bh}$ at the radii of interest. Following \\citet{GM12}, we performed these fits in the radial interval $(r_1 , r_2 )$ such that $N_{\\rm bh}(< r_1) = 5$  and $N_{\\rm bh}(< r_2) = 25$, and assume a density profile of constant power-law index. In  Figure~\\ref{msclu2} we show the (scaled) cumulative number of BHs during the post merger simulations of Section~\\ref{relax}. We find $N_{\\rm bh}(< 0.01{\\rm pc}) \\approx 200$ at $t= 0.1\\times T_{r_{\\rm infl}}$ and $N_{\\rm bh}(< 0.01{\\rm pc}) \\approx 400$ at $t= 0.3\\times T_{r_{\\rm infl}}$. We can directly compare the numbers so obtained with estimates from quasi-steady state Fokker Plank models of the GC. In their  models, \\citet{HA06} assumed $f=0.01$,   and found $N_{\\rm bh}(r < 0.01 {\\rm pc}) = 150$; \\citet{F06} found similar values. This is similar to the number of BHs in our model at the end of the inspiral simulations but somewhat smaller than the predicted number of BHs at the end of the post merger simulations.  \\begin{figure} \\centering $\\begin{array}{ll} \\includegraphics[width=0.44\\textwidth,angle=0]{Figure20.ps} \\end{array}$ \\caption{Eccentricity evolution of test particles taking different numbers of $10M_\\odot$ perturbers: $N_{0.1}$=100, 500, 1000. The initial values of the orbital elements were:  a=10~mpc, e=0.95, $\\Omega=0$, $\\omega=-\\pi/2$ and $i=0.35\\pi$. The diffusion timescale to  evolve to higher angular momenta, where resonant relaxation becomes more efficient in randomizing orbital eccentricities, decreases the larger the number of BHs in the cusp. } \\label{rr-ev} \\end{figure}   In conclusion, a core in  the density distribution of  stars does not necessarily imply a low density of stellar remnants in the GC. In a merger model for NCs, BHs have a smaller core initially -- as a relic of their pre-merger mass segregation, which was not fully erased by the tidal disruption process of the cluster. After a small fraction of the relaxation time~($\\sim 0.1T_{\\rm r_{\\rm infl}}$), the BHs dynamically relax due to BH-BH interactions and attain a steep central density cusp, while the core in the stellar distribution persists. After this time, the number of BHs inside $0.01~$pc, the radii relevant for EMRIs, is $N_{\\rm bh}(<0.01{\\rm pc})\\sim 100$, a value  comparable to that inferred in steady state models.   \\begin{figure} \\centering $\\begin{array}{ll} \\includegraphics[width=0.33\\textwidth,angle=0]{Figure21a.ps} \\\\ \\includegraphics[width=0.33\\textwidth,angle=0]{Figure21b.ps} \\end{array}$ \\caption{Cumulative eccentricity distributions at $5\\times 10^7$yr and at $10^8$yr of evolution for various values of  the number of $10~M_{\\odot}$ BH perutbers inside $0.1$pc, $N_{0.1}$. The insert panels give the $p$-value of the K-S tests as a function of $N_{0.1}$. Only when $N_{0.1}\\gtrsim 1000$ the eccentricity distribution appears to approach a form that is consistent with observations~($p-{\\rm value} \\gtrsim 0.1$) after $\\sim 100~$Myr of evolution (the S-star lifespan). } \\label{s-stars} \\end{figure}  \\subsection{Constraining the number of dark remnants at the GC using the S-star orbits} Observations of the Galctic center have revealed the existence of $\\sim 20$ young (B-type) stars orbiting  within $\\sim 0.05$ pc of the central black hole, the S-stars~\\citep{Ghez08,Gil}. The S-stars are thought not to have formed in situ where the strong tidal field of the MBH would prevent star formation~\\citep{Morriss1993}. Mechanisms that invoke the formation of the S-stars farther from the MBH, and subsequent  rapid migration to their current location  are often invoked  in order to  explain their existence. The current paradigm  is that the S-stars  result from the  capture of individual stars by the tidal disruption of binaries following a  close encounter to the central black hole~\\citep{Hills1988}. This formation model can successfully  accounts for the number of  observed S-stars~\\citep{Perets2007}, but it results into an orbital distribution which is largely biased toward high eccentricities, and therefore it is inconsistent with the quasi-thermal eccentricity distribution of the S-star orbits~\\citep[e.g.,][]{ant+10}. Post-migration dynamical evolution due to gravitational perturbations from a field population of BHs has been invoked in order to bring the predicted orbital distributions more in line with observations~\\citep{Perets2009,Madigan2011,adrian+14}. It has been recently shown that such a scenario could indeed lead to an orbital distribution which reproduces the random and eccentric character of the observed orbits~\\citep{AM13}.  The timescale over which the S-star orbits evolve and approach the observed distribution depends on the number of objects, i.e. BHs, in the stellar cusp, being the shorter the larger the number  of BHs~\\citep{AM13}. Accordingly, we  determine  a lower limit to the number of BHs in the GC by following the dynamical evolution of the S-stars for different values of $N_{\\rm bh}$, and requiring the orbits to approach the observed distribution in  $\\sim100~$Myr, the  main-sequence life time of a B-type star.  We ignore  two-body relaxation since the timescale of interest ($\\sim100~$Myr) is short compared with two-body (non-resonant) relaxation times near the center of the Milky Way~\\citep{M:10}. Following \\citet{AM13} we assume that the semi-major axis distribution, $N(a)$, is known, and it is given by the observed values of $a$. The initial orbital eccentricities were assigned randomly in the range $0.93\\leq e\\leq0.99$. The orbits were initialized at random times between $0$ and $100$~Myr and followed to a final time of  $100$Myr. This setup approximately  reproduces the initial conditions expected in a binary disruption scenario for the formation of the S-stars~\\citep{zhang+13}. For each S-star (i.e., for each value of a), 100  integrations were carried out using the same Hamiltonian model adopted in \\citet{MAMW2011} and \\citet{AM13}. Briefly, the Hamiltonian model accounts for the torques due to finite-$N$ asymmetries in the field-star distribution (resonant-relaxation) and 1-PN terms  which result in the (Schwarzschild) precession of the argument of periastron. The direction of the torquing field was changed smoothly with time and was randomized in a time of order the precession time of the field population as described in Section VB of~\\citet[][]{MAMW2011}.  We consider a field population  of $10~M_\\odot$ BHs,which we distributed according to a power law density profile $ \\rho \\sim  r ^{-2}$. The number of BHs at radii less than $r$ is \\begin{equation} N_{\\rm bh}(<r)=N_{0.1} \\left( r/0.1~{\\rm pc}\\right), \\end{equation} where $N_{0.1} $ is the number of BHs within a radius of $0.1$pc; we take $N_{0.1}=(100,~250,~500,~700,~1000)$, and for each of these values we compared the results of our simulations to observations  by performing a two-sample Kolmogorov-Smirnov~(K-S) test on the eccentricity distributions.  In any of our cusp models,  most orbits have orbital eccentricities that lie initially inside (i.e., at higher eccentricity than) the Schwarzschild barrier -- the locus of points in the ($E,L$) plane where resonant relaxation is suppressed by fast relativistic precession of a test particle orbit~\\citep{MAMW2011}. As shown in \\citet{AM13},  the Schwarzshild barrier behaves as a one-way permeable membrane, which is penetrable  by orbits that approach it from higher eccentricities, while it is a hard barrier for orbits that approach it from above. It can be shown that the time to diffuse above the Schwarzshild barrier, towards higher angular momenta,  scales approximately  with the number of perturbers as $T_{d}\\sim N^{-3/2}_{\\rm bh}(<a)$ or as $\\sim N^{-2}_{\\rm bh}(<a)$ if the choice for the coherence time is the ``vector'' resonant relaxation time or the mass precession time respectively -- i.e., the test particles  diffuse faster toward larger angular momenta, above the Schwarzschild barrier, the larger the number of background perturbers~(Figure~\\ref{rr-ev}).  Figure~\\ref{s-stars} shows the cumulative eccentricity distributions at $5\\times 10^7$yr and at $10^8$yr of evolution for various values of  $N_{0.1}$. As the number of BHs in the cusp increases the diffusion  timescale decreases, allowing the distribution to more rapidly approach a quasi-thermal form. The insert panels give the $p$-value of the K-S tests as a function of $N_{0.1}$. Only for $N_{0.1}\\gtrsim 1000$ after $\\sim 100~$Myr the eccentricity distribution appears to have approached a form that is consistent with observations~($p-{\\rm value} \\gtrsim 0.1$). These results imply that in order to explain the quasi-thermal character of the S-star orbital eccentricities, the number of BHs within $\\sim 0.1~$pc of Sgr~A*  has to be $\\gtrsim1000$. Interestingly, this number is in agreement  with the predictions of our globular cluster merger model (Figure~\\ref{msclu}), but appears to be much larger than the number of BHs implied by the dynamical models discussed  in Section~(\\ref{sec-mass-seg}) which, for $f_{\\rm bh}=10^{-3}$, give $N_{0.1} \\approx 100$ after $10~$Gyr of evolution.  Interestingly, a globular cluster merger model for the Milky Way NC can potentially reconcile models which require high central densities of BHs in order to explain the orbits of the S-stars, with the  ``missing cusp'' problem of the giant star population."
    },
    "1402/1402.0004.txt": {
        "abstract": "{%context } {% aims \\longchange{The geometry of the obscuring ``torus'' in AGN is known to be complex in local sources. At $z=0.5-3$, where most accretion occurs in obscured sources, simple models of line-of-sight (LOS) obscuration are usually assumed. A range of important analyses, such as luminosity functions, accretion density, co-evolution of AGN with host galaxies, rely on results based on these models. We probe the geometry of the torus to find the best model for high-redshift ($z>0.5$) AGN.}{Active Galactic Nuclei are known to have complex X-ray spectra that depend on both the properties of the accreting supermassive black hole (e.g. mass, accretion rate) and the distribution of obscuring material in its vicinity (i.e. the ``torus''). Often however, simple and even unphysical models are adopted to represent the X-ray spectra of AGN, which do not capture the complexity and diversity of the observations. In the case of blank field surveys in particular, this should have an impact on e.g. the determination of the AGN luminosity function, the inferred accretion history of the Universe and also on our understanding of the relation between AGN and their host galaxies.} }{% methods \\longchange{We develop a Bayesian framework for model comparison and parameter estimation of X-ray spectra. We also propagate the uncertainty associated with photometric redshifts estimates. Ten physically motivated models of torus geometries are compared using the deepest X-ray field to date, the Chandra Deep Field South (CDFS) of 4Ms exposure time.}{We develop a Bayesian framework for model comparison and parameter estimation of X-ray spectra. We take into account uncertainties associated with both the Poisson nature of X-ray data and the determination of source redshift using photometric methods. We also demonstrate how Bayesian model comparison can be used to select among ten different physically motivated X-ray spectral models the one that provides a better representation of the observations. This methodology is applied to X-ray AGN in the 4~Ms Chandra Deep Field South. } }{% results \\longchange{For the $\\sim350$ AGN spectra considered, our analysis identifies four important components to represent the variety of AGN spectra: (1) An intrinsic power law, (2) an obscurer which reprocesses the radiation due to photo-electric absorption, Compton scattering and Fe-K fluorescence, (3) an unabsorbed power law associated with Thomson scattering off ionised clouds, and (4) Compton reflection, most noticeable from a stronger-than-expected Fe-K$\\alpha$ line. Simpler models, such as a photo-electrically absorbed power law with a Thomson scattering component, are ruled out with decisive evidence ($Z>100$). We also note that not considering the Thomson scattering component underestimates the inferred column density $N_{H}$. Regarding the geometry of the obscurer, we find strong evidence against both a completely closed, or entirely open, toroidal geometry in favour of an intermediate case.}{For the $\\sim350$ AGN in that field, our analysis identifies four components needed to represent the diversity of the observed X-ray spectra: (1) an intrinsic power law, (2) a cold obscurer which reprocesses the radiation due to photo-electric absorption, Compton scattering and Fe-K fluorescence, (3) an unabsorbed power law associated with Thomson scattering off ionised clouds, and (4) Compton reflection, most noticeable from a stronger-than-expected Fe-K line. Simpler models, such as a photo-electrically absorbed power law with a Thomson scattering component, are ruled out with decisive evidence ($B>100$). We also find that ignoring the Thomson scattering component results in underestimation of the inferred column density, $N_{H}$, of the obscurer. Regarding the geometry of the obscurer, there is strong evidence against both a completely closed (e.g. sphere), or entirely open (e.g. blob of material along the line of sight), toroidal geometry in favour of an intermediate case.} } {% Conclusions \\longchange{Despite the use of low-count spectra, our methodology is able to draw strong inferences on the geometry of the torus. Simpler models are ruled out in favour of a geometrically extended structure with significant Compton scattering. We confirm the presence of a soft component associated with Thomson scattering off ionised clouds in the opening angle of the torus, which scatters the unobscured intrinsic radiation past the torus. The additional Compton reflection may be a sign of a density gradient in the torus or reflection off the accretion disk. Finally, we release a catalogue of AGN in the CDFS with estimated parameters such as the accretion luminosity in the $2-10\\,\\text{keV}$ band and the column density $N_{H}$.}{Despite the use of low-count spectra, our methodology is able to draw strong inferences on the geometry of the torus. Simpler models are ruled out in favour of a geometrically extended structure with significant Compton scattering. We confirm the presence of a soft component, possibly associated with Thomson scattering off ionised clouds in the opening angle of the torus. The additional Compton reflection required by data over that predicted by toroidal geometry models, may be a sign of a density gradient in the torus or reflection off the accretion disk. Finally, we release a catalogue of AGN in the CDFS with estimated parameters such as the accretion luminosity in the $2-10\\,\\text{keV}$ band and the column density, $N_{H}$, of the obscurer.}} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1402/1402.3674_arXiv.txt": {
        "abstract": "We present the first report on an innovative new project named \"RAD@home\", a citizen-science research collaboratory built on free web-services like Facebook, Google, Skype, NASA Skyview, NED, TGSS etc.. This is the first of its kind in India, a zero-funded, zero-infrastructure, human-resource network to educate and directly involve in research, hundreds of science-educated under-graduate population of India, irrespective of their official employment and home-location with in the country. Professional international collaborators are involved in follow up observation and publication of the objects discovered by the collaboratory. We present here ten newly found candidate episodic radio galaxies, already proposed to GMRT, and ten more interesting cases which includes, bent-lobe radio galaxies located in new Mpc-scale filaments, likely tracing cosmological cluster accretion from the cosmic web. Two new Speca-like rare spiral-host large radio galaxies have also been been reported here. Early analyses from our follow up observations with the Subaru and XMM-Newton telescopes have revealed that Speca is likely a new entry to the cluster and is a fast rotating, extremely massive, star forming disk galaxy. Speca-like massive galaxies with giant radio lobes, are possibly remnants of luminous quasars in the early Universe or of first supermassive black holes with in first masssve galaxies. As discoveries of Speca-like galaxies did not require new data from big telescopes, but free archival radio-optical data, these early results demonstrate the discovery potential of RAD@home and how it can help resource-rich professionals, as well as demonstrate a model of academic-growth for resource-poor people in the underdeveloped regions via Internet. ",
        "introduction": "Following discovery of the spiral-host episodic giant radio galaxy `Speca' \\citep{hota11}, it has been observed with the Subaru telescope in long-slit spectroscopy mode and with the XMM-Newton for investigating the kinematics and cluster-environment it belongs to, respectively. H$\\alpha$ and [N II] emission lines seen from the major axis of the galaxy revealed that it is an extremely fast-rotating (maximum speed $\\sim$370 km s$^{-1}$ near the flat regime), star forming disk galaxy. Though it is the brightest member of the identified cluster, the point-source subtracted diffuse X-ray emission is offset from Speca and co-spatial with the early-type members of the cluster. With the discovery of a second such low-z, fast-rotating, spiral-host, episodic giant radio galaxy (Bagchi et al. 2014 presented in this conference), it is clear that they may be rare but it forms a different class of active galaxies. Presence of large number of starbursting, molecular gas-rich (with double-horn profile) radio galaxies at high-z and presence of luminous quasars with billion solar mass supermassive black holes of early Universe suggest their remnants to be found at the low-z Universe. We hypothecise that Speca-like spiral-host, massive, giant radio galaxies are the possible remnants of these first massive galaxies with supermassive black holes, who have avoided major mergers (with big spirals and into clusters) and have become radio loud only recently. So, finding more Speca-like galaxies is like spotting `living fossils' and it is important to investigate their nature, origin with more discoveries for better understanding of black hole, galaxy and cluster evolution. \\\\ ",
        "conclusions": ""
    },
    "1402/1402.0677_arXiv.txt": {
        "abstract": "We present an updated version of the so-called Madau model for the attenuation by the intergalactic neutral hydrogen against the radiation from distant objects. First, we derive a distribution function of the intergalactic absorbers from the latest observational statistics of the Ly$\\alpha$ forest, Lyman limit systems, and damped Ly$\\alpha$ systems. The distribution function excellently reproduces the observed redshift evolutions of the Ly$\\alpha$ depression and of the mean-free-path of the Lyman continuum simultaneously. Then, we derive a set of the analytic functions which describe the mean intergalactic attenuation curve for objects at $z>0.5$. The new model predicts less (or more) Ly$\\alpha$ attenuation for $z\\simeq3$--5 ($z>6$) sources through usual broad-band filters relative to the original Madau model. This may cause a systematic difference in the photometric redshift estimates, which is, however, still as small as about 0.05. Finally, we find a more than 0.5 mag overestimation of the Lyman continuum attenuation in the original Madau model at $z>3$, which causes a significant overcorrection against direct observations of the Lyman continuum of galaxies. ",
        "introduction": "Radiation from cosmological sources is absorbed by neutral hydrogen left in the intergalactic medium (IGM) even after the cosmic reionization \\citep[e.g.][]{gun65}. This intergalactic neutral hydrogen probably traces `cosmic web' produced by the gravity of the dark matter \\citep[e.g.][]{rau98}. Along an observer's line-of-sight piercing the cosmic web, there appear to be numerous discrete systems composed of the intergalactic neutral hydrogen producing a number of absorption lines in the spectra of distant sources. These systems are divided into the Ly$\\alpha$ forest (LAF; $\\log_{10}(N_{\\rm HI}/{\\rm cm}^{-2})<17.2$), Lyman limit systems (LLSs; $17.2\\leq\\log_{10}(N_{\\rm HI}/{\\rm cm}^{-2})<20.3$) and damped Ly$\\alpha$ systems (DLAs; $\\log_{10}(N_{\\rm HI}/{\\rm cm}^{-2})\\geq20.3$), depending on the column density of the neutral hydrogen along the line-of-sight \\citep[e.g.][]{rau98}.  The intergalactic absorption is routinely found in the spectra of objects at a cosmological distance and this feature is utilized as a tool to select high-$z$ objects only by photometric data: the so-called drop-out technique \\citep[e.g.][]{ste95,mad96}. To put it another way, we must always correct the spectra of cosmological sources for this absorption in order to know the intrinsic ones. Therefore, an accurate model of this absorption is quite useful as a standard tool for observational cosmology.  After several models for this purpose were presented \\citep[e.g.,][]{mol90,zuo93,yos94}, \\cite{mad95} (hereafter M95) appeared and became the most popular because of the convenient analytic functions. However, the heart of the model, i.e. the statistics of the LAF, LLSs, and DLAs, has been updated largely by observations in the last two decades since M95. In fact, there are several papers adopting such updated statistics \\citep{ber99,mei06,tep08,ino08}. Nevertheless, people still adhere to M95, except for a few innovative authors \\citep[e.g.][]{har11,ove13}. This adherence may be due to the simplicity and convenience of the analytic functions in M95. Here, in this paper, we intend to present a user-friendly analytic function conforming to updated statistics.  In the next section, we introduce the heart of the modeling: the distribution function of the intergalactic absorbers derived from the latest observational data of the LAF, LLSs, and DLAs. Then, we show the updated mean transmission function and compare it with the latest observations of the Ly$\\alpha$ transmission and the mean-free-path of Lyman limit photons in section 3. In section 4, we present new analytic formulae of the intergalactic attenuation. Finally, we quantify the difference of the attenuation magnitudes through some broad-band filters between the M95 model and ours and discuss the effect on the drop-out technique and the photometric redshift estimation in section 5. We do not need to assume any specific cosmological model in this paper, except for sections 2.1 and 3.2, where we assume $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Therefore, the analytic functions presented in section 4 can be used directly in any cosmological models. ",
        "conclusions": "We here compare the attenuation magnitudes through some broad-band filters for the three models discussed in this paper, quantify the difference, and discuss the effect on the drop-out technique and the photometric redshift (hereafter photo-$z$) estimation.  Suppose the flux density observed through the IGM at the wavelength $\\lambda$ is expressed as $F_\\lambda^{\\rm obs}=F_\\lambda^{\\rm em}T_\\lambda^{\\rm IGM}$, where $F_\\lambda^{\\rm em}$ is the emitted flux density at the proximity of a cosmological object and $T_\\lambda^{\\rm IGM}$ is the IGM transmission. We assume a simple power law spectrum for $F_\\lambda^{\\rm em}$ with a rest-frame ultra-violet index $\\beta_{\\rm UV}$: $F_\\lambda^{\\rm em} \\propto \\lambda^{\\beta_{\\rm UV}}$. A band magnitude for using a photon-counting detector like CCDs is defined by $m=-2.5\\log_{10} F + C_0$ with $F=\\int (F_\\nu t_\\nu/h\\nu) d\\nu / \\int (t_\\nu/h\\nu) d\\nu=\\int F_\\lambda t_\\lambda (\\lambda/c) d\\lambda / \\int (t_\\lambda / \\lambda) d\\lambda$, where $t_\\nu=t_\\lambda$ is the total (including the filter, detector, telescope and instrument optics, and atmosphere) efficiency of the band, and $C_0$ is the magnitude zero point. Thus, we can express the IGM attenuation through a band filter as \\begin{equation} \\Delta m_{\\rm IGM} = -2.5 \\log_{10} \\left( \\frac{\\int F_\\lambda^{\\rm obs} t_\\lambda \\lambda d\\lambda} {\\int F_\\lambda^{\\rm em} t_\\lambda \\lambda d\\lambda} \\right) = -2.5 \\log_{10} \\left( \\frac{\\int \\lambda^{\\beta_{\\rm UV}+1} T_\\lambda^{\\rm IGM} t_\\lambda d\\lambda}{\\int \\lambda^{\\beta_{\\rm UV}+1} t_\\lambda d\\lambda} \\right)\\,. \\end{equation}  Figure~8 shows the IGM attenuation magnitudes through 6 broad-band filters as a function of the source redshift in the case of $\\beta_{\\rm UV}=-2.0$, a flat continuum in $F_\\nu$ unit usually observed in high-$z$ star-forming galaxies \\citep[e.g.,][]{sha03}. We note here that the variation of $\\beta_{\\rm UV}$ from $-3$ to 0 \\citep[e.g.,][]{bou09} has a negligible effect on the attenuation magnitudes. The solid, dotted, and dashed lines are the models of this work, II08, and M95, respectively. The attenuation magnitudes shown in the figure are determined mainly by Ly$\\alpha$ and Ly$\\beta$ absorptions. Then, the difference seems to be small as expected from the small difference of the mean transmission curves among the three models shown in Figure~4. In fact, however, the vertical difference at a fixed source redshift reaches more than 1 mag between this work and the M95 model, while the horizontal difference is as small as about $<0.2$, except for the deviation of the M95 model at $z_{\\rm S}>5.5$ owing to the lack of the rapid evolution of the Ly$\\alpha$ optical depth included in the other two models. The thin (coloured) solid lines are the results from the analytic formulae for the new model presented in the previous section. We find an excellent agreement with the numerical integrations.  \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{fig8.eps} \\end{center} \\caption{IGM attenuation magnitude through broad-band filters, the Canada-France-Hawaii Telescope/Mega-cam $u^*$, and the Subaru/Hyper Suprime-Cam $g$, $r$, $i$, $z$ and $y$, as a function of the source redshift. The solid, dotted, and dashed lines are the models of this work, Inoue \\& Iwata (2008), and Madau (1995), respectively. The thin (coloured) solid lines are the cases using the analytic formulae for this work. The object spectrum with the ultra-violet spectral index $\\beta_{\\rm UV}=-2.0$ is assumed.} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{fig9.eps} \\end{center} \\caption{Difference of the means of the photometric redshift estimations assuming the IGM model of this work relative to those assuming the Madau (1995) model. The sample is 427 galaxies with spectroscopic redshift larger than 2 in the GOODS-S field and is divided into bins 20 objects each. Along the horizontal axis, the points and error-bars show the mean and the standard deviation of the spectroscopic redshifts in each bin. See the text for the vertical error-bars.} \\end{figure}  Although the horizontal difference at a certain amount of the attenuation magnitudes among the three models is small, there is a difference which would affect the drop-out technique and the photo-$z$ estimation. The drop-out threshold is usually $\\Delta m_{\\rm IGM}\\simeq1$ mag. The source redshift reaching the threshold is different from the models. For example, the redshifts in the M95 model are about 0.2 smaller than those of this work at $z_{\\rm S}\\simeq3$--4 but are about 0.1 larger at $z_{\\rm S}\\simeq6$. These difference would result in systematically lower or higher photo-$z$ solutions with the M95 model than with the new model of this paper. To check this expectation, we ran a photo-$z$ code developed by \\cite{tan13a,tan13b} adopting two IGM models of this paper and M95. The sample is the galaxies with spectroscopic redshifts and photometry of VLT/VIMOS $U$, HST/ACS F435W, F606W, F775W, F814W, F850LP, HST/WFC3 F105W, F125W, F160W, VLT/ISAAC $K_{\\rm s}$, Spitzer/IRAC {\\it Ch1} and {\\it Ch2} in the GOODS-S field \\citep{guo13}. We collected spectroscopic redshifts from the literature \\citep{lef05,mig05,van08,pop09,bal10} and cross-matched with the photometric objects within 1 arcsec. We use secure redshifts only in the analysis here. The photo-$z$ code assumes the stellar population synthesis model by \\cite{bru03} with solar and sub-solar metallicity models ($Z=0.02$, 0.008, and 0.004), the Chabrier initial mass function between 0.1 and 100 $\\rm M_\\odot$ \\citep{cha03}, exponentially declining star formation history, the Calzetti attenuation law \\citep{cal00}, and the emission line model by \\cite{ino11a} with the Lyman continuum escape fraction of zero. The Ly$\\alpha$ emission line is reduced by a factor of 0.1 to account for the attenuation through the interstellar medium of galaxies. The metallicity, age, exponential time-scale of the history, attenuation amount, and redshift are free parameters determined by a $\\chi^2$ minimization technique. We compare the photo-$z$s for the two IGM models in Figure~9. We divided the sample galaxies into bins 20 objects each and calculated the difference of the means of the photo-$z$s in each bin. The vertical error-bars are estimated by $\\sqrt{\\sum_i(\\sigma_{\\langle z_{{\\rm ph},i}\\rangle}^2/n+\\delta_{z_{{\\rm ph},i}}^2/n)}$, where $i$ indicates the two IGM models (this work and M95), $\\sigma_{\\langle z_{{\\rm ph},i}\\rangle}$ is the standard deviation of the photo-$z$s in each bin, $\\delta_{z_{{\\rm ph},i}}$ is the mean of photo-$z$ uncertainties of the sample galaxies in each bin, and $n=20$ is the number of the sample galaxies in each bin. The first term is the standard error of the mean and the second term is the error in the mean propagated from the uncertainty of the individual photo-$z$. As found in Figure~9, the difference of the means of photo-$z$s is too small to be detected in the sample adopted, while we may find the expected trend of photo-$z$s for the new IGM model larger (or smaller) than those for the M95 model at $z\\simeq3$--5 ($z>5.5$). The marginal difference of $\\approx0.05$ at $z\\simeq3.5$ is much smaller than that expected from Figure~8. This is probably because we used not only the drop-out band but also all bands available in the photo-$z$ estimation. As a result, the drop-out feature has a lower weight in the photo-$z$ determination. However, all the available bands should be used in order to constrain intrinsic shapes of the spectral energy distribution below Ly$\\alpha$ to characterize the IGM effect on photo-$z$. We would detect the IGM model difference securely if we had a ten times larger number of the sample galaxies at $z>3$.  \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{fig10.eps} \\end{center} \\caption{IGM attenuation magnitude at the rest-frame 880 \\AA\\ as a function of the source redshift. The solid, dotted, and dashed lines are the models of this work, Inoue \\& Iwata (2008), and Madau (1995), respectively.} \\end{figure}  Finally, we examine the mean IGM attenuation magnitude at a Lyman continuum wavelength 880 \\AA\\ in the source rest-frame as a function of the source redshift in Figure~10. This is motivated by studies for determining an important parameter controlling the cosmic reionization, the Lyman continuum escape fraction of galaxies \\citep[e.g.,][]{ino05,iwa09}. In these studies, we need to correct the IGM attenuation against the observed Lyman continuum of galaxies. As found in Figure~10, the difference among the three models discussed in this paper is significant; the new model predicts the least attenuation which is 0.5--1 mag smaller than the M95 model at $z_{\\rm S}=3$--4. This is consistent with those found in Figures~4 and 6. Note that this less attenuation against the Lyman continuum comes from the recent updates of the occurrence rate of LLSs discussed in section 2.2 and the measurements of the mean-free-path discussed in section 3.2. Therefore, using the M95 model causes a significant overcorrection of the observed Lyman continuum and results in an overestimation of the escape fraction. On the other hand, the Lyman continuum absorption is mainly caused by LLSs which are relatively rare to have on a line-of-sight. As a result, a large fluctuation of the attenuation amount among many lines-of-sight is expected. Therefore, a Monte-Carlo simulation is required to model the stochasticity as done in II08. This point would be investigated in our next work."
    },
    "1402/1402.2672_arXiv.txt": {
        "abstract": "We present a systematic search for molecular hydrogen absorption systems at high redshift in quasar spectra from the Sloan Digital Sky Survey (SDSS) II Data Release 7 and SDSS-III Data Release 9. We have selected candidates using a modified profile fitting technique taking into account that the Ly$\\alpha$ forest can effectively mimic H$_2$ absorption systems at the resolution of SDSS data. To estimate the confidence level of the detections, we use two methods: a Monte-Carlo sampling and an analysis of control samples. The analysis of control samples allows us to define regions of the spectral quality parameter space where H$_2$ absorption systems can be confidently identified. We find that H2 absorption systems with column densities $\\log {\\rm N_{H_2}} > 19$ can be detected in only less than 3\\% of SDSS quasar spectra. We estimate the upper limit on the detection rate of saturated H$_2$ absorption systems ($\\log {\\rm N_{H_2}} > 19$) in Damped Ly-$\\alpha$ (DLA) systems to be about 7\\%. We provide a sample of 23 confident H$_2$ absorption system candidates that would be interesting to follow up with high resolution spectrographs. There is a 1$\\sigma$\\,\\,$r-i$ color excess and non-significant $A_{\\rm V}$ extinction excess in quasar spectra with an H$_2$ candidate compared to standard DLA-bearing quasar spectra. The equivalent widths (EWs) of C\\,{\\sc ii}, Si\\,{\\sc ii} and Al\\,{\\sc iii} (but not Fe\\,{\\sc ii}) absorptions associated with H$_2$ candidate DLAs are larger compared to standard DLAs. This is probably related to a larger spread in velocity of the absorption lines in the H$_2$ bearing sample. ",
        "introduction": "\\label{introduction} \\noindent The Sloan Digital Sky Survey (SDSS; \\citealt{York2000}) is one of the largest optical surveys of modern astrophysics. One of the major goals of this survey is to study  large scale structures in the nearby Universe $z<0.7$, from the spatial distribution of galaxies \\citep{York2000}. The survey targets millions of objects (galaxies, stars and quasars) by imaging and spectroscopy in the optical wavelength band. The recently published ninth data release 9 (DR9) \\citep{Ahn2012} contains over 80\\,000 spectra of high redshift quasars \\citep{Paris2012}.  Absorption lines in quasar spectra allow one to study the intergalactic and interstellar matter located along the line of sight to the quasar. Quasars are routinely detected up to $z\\sim6$, which corresponds to more than 12 Gyr ago. The Baryon Oscillation Spectroscopic Survey (BOSS; Schlegel et al. 2007; Dawson et al. 2012), one of the main projects of  SDSS-III (the third generation of SDSS) is primarily focused on the analysis of the spatial distribution of luminous red galaxies (LRGs) and absorptions in the Ly$\\alpha$ forest. The later enables to determine the baryon acoustic oscillation (BAO) scale at z\\,$\\sim 2.5$ that corresponds to an epoch before dark energy dominates the expansion of the Universe \\citep{Busca2013}. In addition to this primary goal, BOSS spectra will allow one to study of broad absorption line (BAL) systems arising in the vicinity of Active Galaxy Nucleus (AGN), as well as intervening metal absorption line systems (such as C~{\\sc iv} and Mg~{\\sc ii}) associated with clouds located into the halo of intervening galaxies \\citep{Quider2011, Zhu2013}.  Damped Lyman Alpha systems (DLAs) are identified by a broad H~{\\sc i} Ly-$\\alpha$ absorption line with prominent saturated Lorentz wings. Statistical analysis shows that DLA systems are the main reservoir of neutral gas at high redshifts \\citep{Prochaska2009, Noterdaeme2009}. It is believed that these systems are associated with disks of galaxies or their close vicinity with impact parameter less than 20 kpc \\citep{Fynbo2011, Krogager2012}. Numerous species are seen in DLAs (e.g. C\\,{\\sc i} to C\\,{\\sc iv} and Si\\,{\\sc i} to  Si\\,{\\sc iv}) and the typical velocity spread of metal absorption lines is about 100--500 km/s. All this makes DLA systems relatively easy to detect at intermediate resolution. About 12,000 DLA/sub-DLA candidates with log N(H\\,{\\sc i})$>=$20 were detected in SDSS-DR9 \\citep{Noterdaeme2012}. It has been shown that a small fraction of DLAs hosts molecular hydrogen \\citep{Noterdaeme2008} which corresponds to diffuse and translucent neutral clouds embedded in warm neutral interstellar medium. In some cases HD molecules are detected \\citep{Varshalovich2001, Noterdaeme2008b, Balashev2010} as well as CO molecules \\citep{Srianand2008, Noterdaeme2011}. Observations of molecular hydrogen clouds at the high redshift Universe provide a unique opportunity to study several issues. Molecular hydrogen is believed to be an indicator of the cold phase of the interstellar medium which is the raw material for the star-formation -- it was found that regions with high H$_2$ abundance are correlated with the star-formation regions \\citep{Krumholz2012}. The measurement of relative abundances of different H$_2$ rotational levels allows one to determine the physical conditions in this cold interstellar medium --  kinetic temperature, UV radiation field and possibly number density. There are several cosmological problems related to high redshift molecular hydrogen absorption systems: (i) the detection of HD gives the possibility of a complementary approach to the determination of the primordial deuterium abundance \\citep{Ivanchik2010,Balashev2010}; (ii) constraints on the possible variation of the proton to electron mass ratio can be obtained (e.g. \\citealt{Thompson1975, Ivanchik2005, Wendt2012, Rahmani2013}); (iii) the CMBR temperature can be measured at high redshift from the populations of the CO rotational levels (e.g. \\citealt{Noterdaeme2011}) and  C\\,{\\sc i} fine-structure levels (e.g. \\citealt{Songaila1994, Srianand2000}).  Since the first detection by \\citet{Levshakov1985}, about 20 H$_2$ absorption systems have been detected at high redshifts ($z>1$, see \\citealt{Ge1997, Ledoux2002, Ledoux2003, Reimers2003, Cui2005, Noterdaeme2008, Srianand2008, Jorgenson2009, Jorgenson2010, Malec2010, Noterdaeme2010, Fynbo2011, Guimaraes2012}). Additionally, two systems were detected in spectra of GRB afterglows \\citep{Prochaska2009b, Kruhler2013}, one system was detected at intermediate redshift $z\\sim 0.5$ \\citep{Crighton2013}, and four systems \\citep{Noterdaeme2009, Noterdaeme2011} were detected by means of CO molecules (H$_2$ and HD transitions in these systems are redshifted out of the observed wavelength range). Most of the detection of molecular hydrogen absorption systems have been performed with high signal to noise ratio (S/N) and high resolution spectra with typically 3 hours exposures on a 8~m class telescope. On the other hand, there are about 80,000 high redshift quasar spectra obtained in the course of SDSS to be search for H$_2$. The main disadvantages of SDSS spectra for detection of molecular hydrogen systems are that they have intermediate spectral resolution $R$$\\sim$2000 and usually low S/N ($<$4 over the wavelength range where H$_2$ absorption lines are to be found).  In this paper we show however that it is possible to find H$_2$-bearing systems in SDSS spectra provided that the column density is large enough. Although such systems are  rare, the large number of SDSS quasar spectra makes their detection possible. We provide a list of H$_2$ candidates for follow-up with high resolution observations. The paper is organized as follows. In Section~2 we give a brief description of the data. Section~3 describes a searching criterion to be used. We define a quantitative estimate of the confidence level or false identification probability in Section~4. The final list of the H$_2$ system candidates is given in Section~5 and we investigate some properties of this sample in Section~6 before conclusions are drawn in Section~7. ",
        "conclusions": "\\label{conclusion} \\noindent  We have performed a systematic search for H$_2$ absorption in DLAs detected  towards quasars in the Sloan Digital Sky Survey releases DR7 and DR9. We have developed and used fully automatic procedures based on the profile fitting technique with modified $\\chi^2$ function. The main difficulty of the search is the presence of the Ly-$\\alpha$ forest which is dense at redshifts under consideration ($z>2.2$) and can effectively mimic H$_2$ absorption lines at the resolution and S/N of SDSS spectra. We carefully checked our searching procedure and found that it yields a completeness of $>99\\%$ for $\\log$N$_{\\rm H2}$ $>$18.5. However the number of false detections is rather high.  For each of the candidates, the probability of false identification was calculated using two techniques. The first technique employs Monte-Carlo simulations to estimate the probability of an accidental fit of an absorption system in the Ly$\\alpha$ forest. In this technique we used repeated random shifts for each H$_2$ absorption line in the particular spectrum. The advantage of this technique is to give an estimate of the false identification probability in the particular spectrum (i.e. the particular realization of the Ly$\\alpha$ forest). The second technique uses a control sample to estimate the false identification probability. Control sample consists of the non-BAL non-DLA quasars from DR9 SDSS catalog which guarantees the absence of H$_2$ absorption systems. The false identification probability can be calculated as the identification rate of H$_2$ absorption systems in the control sample. We additionally have performed simulations of mock SDSS quasar spectra which were used to check the procedures and to increase statistics. We have found that the false identification probability depends on the H$_2$ column density and spectral properties, mainly on S/N, $z_{\\rm QSO}$, and the number  of H$_2$ absorption lines present in the spectrum. This method allows us to estimate the detection limit of H$_2$ absorption systems in SDSS spectra, i.e. the spectral properties for which the false identification probability is low and therefore the identification of the H$_2$ absorption system is robust. We found that in SDSS data a reasonable detection limit for H$_2$ column density, $\\log$N$_{\\rm H_2}$, is higher than 19. We also derived that the upper limit on the fraction of DLAs with H$_2$ systems is statistically equal to 7 and 3\\% for, respectively, $\\log$N$_{\\rm H_2}>$19 and 19.5. These are as upper limits because we tend to overestimate H$_2$ column densities. The false identification probabilities calculated by the Monte-Carlo sampling and control sample techniques agree with each other very well.  We have selected 23 candidates of H$_2$ absorption systems with high confidence level which are promising candidates for follow up high resolution studies. Taking into account the derived upper limit on the fraction of DLAs with saturated H$_2$ lines $\\sim 7 \\%$ it leads us to the conclusion that only less than 3\\% of SDSS spectra are suitable for confident detection of H$_2$ absorption systems.  We studied the properties of these candidates, namely, color excess, extinction and equivalent widths of metal lines. There is a 1$\\sigma$ $r-i$ color excess and no significant $A_V$ extinction in quasar spectra with an H2 candidate compared to standard DLA bearing quasar spectra. We find larger C\\,{\\sc ii}, Si\\,{\\sc ii}, and Al\\,{\\sc iii} equivalent widths (EW) in the H$_2$ candidate sample over the DLA sample but no significant difference for Fe\\,{\\sc ii} EWs. This is probably related to a larger spread in velocity of the absorption lines in the H$_2$ bearing sample and therefore possibly larger metallicities.  The selected candidates would increase by a factor of two the number of known H$_2$ absorption systems at high redshift. It is important to note, that we can confidently identify only saturated H$_2$ systems with high column densities, $\\log$N$_{\\rm H_2} > 19$. To date there are only a few such systems detected at high redshift. The selected candidates will undoubtedly allow us to gather a large sample of HD detections. The relative abundance of HD/H$_2$ molecules which provides important clues on the chemistry and star-formation history and also can be used to estimate the isotopic D/H ratio and consequently $\\Omega_b$ -- the density of baryonic matter in the Universe \\citep{Ivanchik2010, Balashev2010}. In addition these systems are unique objects to search for CO molecule \\citep{Srianand2008}. This molecule is suitable to estimate the physical conditions in interstellar clouds and the CMBR temperature at high redshift \\citep{Noterdaeme2010, Noterdaeme2011}. We have selected four candidates with column densities $\\log$N$_{\\rm H_2} > 20$. Such systems have never been observed at high redshifts towards QSO sightlines and would provide an exclusive opportunity to study translucent clouds in the interstellar medium at high redshift.  \\vspace{2mm}{\\footnotesize {\\rm Acknowledgments.} We are very grateful the anonymous referee for the detailed and careful reading our manuscript and many useful comments. This work was partially supported by a State Program ``Leading Scientific Schools of Russian Federation'' (grant NSh-294.2014.2). SB and SK partially supported by the RF Presedent Programme (grant MK-4861.2013.2)."
    },
    "1402/1402.0488_arXiv.txt": {
        "abstract": "Neutron stars are thought to be excellent laboratories for determining the equation of state (EoS) of cold dense matter. Their strong gravity suggests that they can also be used to constrain gravity models. The two observables of neutron stars---mass and radius (M-R)---both depend on the choice of EoS and relativistic gravity, meaning that neutron stars cannot be simultaneously good laboratories for both of these questions. A measurement of mass and/or radius would constrain the less well known physics input. The most common assumption---namely, that M-R measurements can be used to constrain EoS---presumes general relativity (GR) is the ultimate model of gravity in the classical regime. We calculate the radial profile of compactness and curvature (square root of the full contraction of the Weyl tensor) within a neutron star and determine the domain not probed by the Solar System tests of GR. We find that, except for a tiny sphere of radius less than a millimeter at the center, the curvature is several orders of magnitude above the values present in Solar System tests. The compactness is beyond the solar surface value for $r>10$ m, and increases by 5 orders of magnitude towards the surface. With the density being only an order of magnitude higher than that probed by nuclear scattering experiments, our results suggest that the employment of GR as the theory of gravity describing the hydrostatic equilibrium of the neutron stars is a rather remarkable extrapolation from the regime of tested validity, as opposed to that of EoS models. Our larger ignorance of gravity within neutron stars suggests that a measurement of mass and/or radius constrains gravity rather than the EoS, and given that the EoS has yet to be determined by nucleon scattering experiments, M-R measurements cannot tightly constrain the gravity models either. Near the surface the curvature and compactness attain their largest values,  while the EoS in this region is fairly well known. This renders the crust as the best site to look for deviations from GR. \\\\ \\vspace{1cm} \\doi{10.1103/PhysRevD.89.063003} \\hspace{2.7cm} PACS numbers: 97.60.Jd, 04.40.Dg, 04.80.Cc ",
        "introduction": "\\label{intro}  \\thispagestyle{fancyplain}  Neutron stars with their diversity of astrophysical manifestations provide many challenging problems that make them interesting in their own right. Two aspects of neutron stars, however, make these objects especially important for fundamental physics: (i) the high densities prevailing at the core of neutron stars make them possible laboratories for determining the equation of state (EoS) of the ultimate ground state of cold catalyzed dense matter (e.g., Refs. \\citep []{guv+1820,guv+1608,guv+1745,oze+10,guv13,ste+10,oze+12,neu+12,sul+11}; see Refs. \\citep[]{lat07,lat12,han+07} for reviews); (ii) the strong gravitational fields of neutron stars allow them to be used as tools for testing or constraining alternatives \\citep[]{har98,sot04,yag+13,hor11,sot12,las+08,pan+11} or modifications \\citep[]{coo+10,ara+11,del+12,odd+23,odd+24} of general relativity (GR) (see Refs. \\citep[]{wil06,wil09,psa08} for reviews). Neutron stars cannot simultaneously be good test beds for both questions (see, e.g., Refs. \\citep[]{wen+11}). The often-mentioned notion that neutron stars are excellent labs to constrain the EoS presumes that GR is the ultimate theory of gravity at the classical level. This seems to be a safe assumption given that the parametrized post-Newtonian (PPN) parameter $\\gamma$ (a measure of curvature) has been tested with a precision of $10^{-5}$ \\citep{ber+03,sha+04} and found to be consistent with the prediction of GR, $\\gamma=1$. As argued in the excellent review of Psaltis \\citep{psa08}, however, these Solar System tests constrain only weak-field gravity; gravity in the strong-field regime is largely unconstrained by observations. The differences between GR and its alternatives/modifications can become prominent in the strong gravitational fields of neutron stars \\citep{ded03}.  The mean density of a typical neutron star is $\\bar{\\rho}=3M_{\\ast}/4\\pi R_{\\ast}^3 = 0.67\\times 10^{15}M_{1.4}R_{6}^{-3}$ g cm$^{-3}$ ($M_{1.4}=M_{\\ast}/1.4M_{\\odot}$ and $R_{6}=R_{\\ast}/10^6\\, {\\rm cm}$). This is  an order of magnitude higher than the densities probed by nucleon scattering experiments \\citep[e.g.][]{dan+02}.  At the surface of the star the compactness $\\eta_{\\ast} \\equiv 2GM_{\\ast}/R_{\\ast}c^2 = 0.416\\,M_{1.4}R_{6}^{-1}$ and ``curvature'' ${\\cal K}_{\\ast} \\equiv 4\\sqrt{3}GM_{\\ast}/R_{\\ast}^3 c^2 = 1.18\\times 10^{-12}\\,M_{1.4}R_{6}^{-3}$~cm$^{-2}$ (see Sec. \\ref{sec-method} for the definition of curvature), the two measures of the gravitational field strength \\citep{psa08}, are 5 and 14 orders of magnitudes larger, respectively, than their counterparts in the Solar System ($\\eta_{\\odot} \\equiv 2GM_{\\odot}/R_{\\odot}c^2 = 4.27\\times 10^{-6}$ and ${\\cal K}_{\\odot}  \\equiv 4\\sqrt{3}GM_{\\odot}/R_{\\odot}^3 c^2 =3.06\\times 10^{-27}$~cm$^{-2}$). This large unconstrained regime in gravity compared to that in the EoS may not immediately imply that a measurement of the mass and/or radius (M-R) of a neutron star  would constrain gravity models rather than the EoS of dense matter: the density of a neutron star is nearly constant at the core and drops by 15 orders of magnitude near the surface meaning that a large fraction of the mass and radius of the neutron star might be determined by the unconstrained EoS regime, while the radially increasing compactness and curvature within the star could be in a well-tested regime for most parts of the star. Such an argument is probably the reason for the persistence of the general consent within the astrophysics community that a sufficiently accurate M-R measurement would constrain the EoS at the core of neutron stars rather than gravity. Yet another reason could be that GR is singled out among its alternatives as being the simplest valid model of gravity, whereas there is no EoS that is equally as prominent. To decide which---gravity or the EoS---is better understood, it is necessary to  quantify where within the star the strength of gravity is beyond the values probed in Solar System tests.   The two future X-ray missions---NASA's Neutron star Interior Composition ExploreR (NICER) \\citep{gen+12} and ESA's the Large Observatory For X-ray Timing (LOFT) \\citep{fer+12}---are expected to measure the mass and radius of neutron stars \\citep{lo+13,psa+13}. The measurement of the mass and/or radius of a neutron star would constrain the less well known physics input more than the other. Whether results from these missions would constrain the EoS or gravity models depends on which of these physics inputs is less well known within neutron stars. The purpose of this work is to quantify the range of unconstrained gravity within neutron stars in the framework of GR. This clarifies how much the M-R relations obtained within the framework of GR extrapolate its usage from its well-tested regime. In Sec. \\ref{sec-method} we present the hydrostatic equilibrium equations and the definitions of the compactness and curvature that we employ. In Sec. \\ref{sec-results} we present our results showing the radial profiles of the compactness and curvature within the neutron star. Finally, in Sec. \\ref{sec-discuss} we discuss our results. We present the derivations of the curvature scalars in an Appendix. ",
        "conclusions": "\\label{sec-discuss}  Our results demonstrate that the strength of gravity within a neutron star is orders of magnitude larger than the gravitational strengths probed in Solar System tests, indicating that the presupposition of the validity of GR in determining the neutron star structure is an ``extravagant'' extrapolation. The EoS, on the other hand, is only an order of magnitude larger than what is being probed in nucleon scattering experiments. This does not only imply that an M-R measurement would constrain gravity models more than EoS models, but also that---given that EoSs produce very different M-R relations---it would not be possible to investigate deviations from GR by any measurement of M-R unless the EoS is fully constrained by Earth-based experiments. The strong gravity of neutron stars, nevertheless, might allow one to constrain modified gravity models in a perturbative approach (see, e.g., Refs. \\citep[]{coo+10,ara+11,del+12}).  Our results are consistent with that of DeDeo and Psaltis \\citep{ded03}, where the authors showed that M-R relations obtained by different models of gravity show a much greater variance than M-R relations obtained with different EoSs. Our results complement those of DeDeo and Psaltis \\citep{ded03} by quantifying the strength of gravity \\emph{within} the neutron stars and clarifies the fact that the domain where gravity is in a not-well-tested regime largely overlaps with the domain where the EoS is not constrained.  Our results for the domain of tested validity obviously depend on the gravity model chosen, i.e., GR. The analysis we provide can be repeated for different viable gravity models with parameters constrained  by Solar System tests. Though such a systematic survey is beyond the scope of this paper, the present constraints on the mass and radius taken roughly as $M_{\\ast}=1-2\\,M_{\\odot}$ and $R=8-15$~km, already lead to a very large compactness and curvature at the surface. These values obtained via phenomena outside the neutron star can be taken to be independent of the gravity model employed because almost all models of gravity yield the same vacuum solutions obtained with GR \\citep{psa+08}. It then seems obvious that in any gravity model the curvature and compactness within the neutron star should increase radially from zero to values that are orders of magnitude larger than those probed in the Solar System tests. Hence any model of gravity when employed for the structure of neutron stars would actually be extrapolated from its domain of tested validity.  A systematic stretching of GR beyond its regime of tested validity  by several orders of magnitude naturally raises the question of whether there is any bias in the historical development of EoSs, i.e., whether they have evolved towards producing reasonable M-R relations when used within the framework of GR. Such a question becomes more significant given that astrophysical measurements \\citep{guv+1820,guv+1608,guv+1745,oze+10,guv13,oze+12,ste+10,gui+13}, except for that of Refs. \\citep{sul+11} and \\citep{neu+12}, yield radii systematically smaller than that implied by the Lead Radius Experiment (PREX) \\citep{abr+12}  within the framework of GR \\citep{fat12}.  GR sets a maximum mass $M_{\\max}$ for neutron stars \\citep{rho74} with a different value for each EoS (see Refs. \\citep[]{cha+13,hae12} for a review). The ``two solar mass'' pulsars PSR J1614$-$2230 \\citep{dem+10} and PSR J0348$+$0432 \\citep{ant+13} suggest that $M_{\\max} > 2 M_{\\odot}$. The statistical analysis of K{\\i}z{\\i}ltan \\emph{et al}.\\ \\citep{kiz+13} further suggests that the distribution of masses of neutron stars in nature are not cut off at the upper end such that stellar evolution could produce a low number of $\\simeq 2.5\\,M_{\\odot}$ neutron stars. Within GR this favors stiff EoSs that to yield maximum allowable masses of neutron stars larger than 2.5 solar masses. It is hard to reconcile this result with the small measured radii that favor soft EoSs within GR. This may signal deviations from GR in strong gravity unless there are systematic errors of order 30\\% in the radius measurements. We are yet unaware of any gravity model that yields 2.5 solar-mass objects with 9-11~km with the present EoSs that is also consistent with the Solar System tests. Recently, Arapo{\\u g}lu \\emph{et al}. \\citep{ara+11} showed within an $f(R)=R+\\alpha R^2$ gravity model that the maximum mass of neutron stars becomes larger for smaller values of $\\alpha$, while the radius remains nearly constant. This model of gravity then gives large-mass objects for $\\alpha<0$ even with soft EoSs. The nonrelativistic limit of such a gravity model, however, requires $\\alpha>0$ \\citep{naf10}, rendering solutions with $\\alpha <0$ nonphysical.  Both the compactness and the curvature reach their maximum values near the crust. Considering that the EoS of the crust is fairly well known \\citep{BPS,NV}, we are led to the conclusion that the crust of neutron stars provides the best sites in the Universe for identifying deviations from GR in strong gravity. Such phenomena displayed by the crust are, e.g., the glitch activity of young pulsars \\citep{alp05} and modulation of the cooling of neutron stars \\citep{yak04}. One may then expect deviations from GR to be identified from, e.g.,\\ the glitch activity of pulsars given that GR is the least constrained physics input in the description of the phenomena by 14 orders of magnitude in curvature and 5 orders of magnitude in compactness, according to Fig.~\\ref{fig-curv_log}. Such a deviation from GR may already have been observed! The recent inference \\citep{and+12} that the moment of inertia of the crust of neutron stars is not sufficient to understand some of the glitch activity through the highly successful vortex creep model \\citep{alp+84,alp+93} may signal a  deviation from GR in strong gravity rather than a requirement for the modification of the model itself. An investigation of the effect of scalar fields \\citep{dam93} on the moment of inertia of neutron stars is left for a future work.   \\begin{center} {\\bf ACKNOWLEDGEMENTS} \\end{center}  \\noindent We acknowledge support from the scientific and technological council of Turkey (T\\\"{U}B\\.{I}TAK) with the project number 112T105. This work was supported by the COST Action  MP1304."
    },
    "1402/1402.0513_arXiv.txt": {
        "abstract": "We investigated the impacts of magnetorotational instability (MRI) on the dynamics of weakly magnetized, rapidly rotating core-collapse by conducting high resolution MHD simulations in axisymmetry with simplified neutrino transfer. We found that an initially sub-magnetar class magnetic field is drastically amplified by MRI and substantially affects the dynamics thereafter. Although the magnetic pressure is not strong enough to eject matter, the amplified magnetic field efficiently transfers angular momentum from higher to lower latitudes, which causes the expansion of the heating region at low latitudes due to the extra centrifugal force. This then enhance the efficiency of neutrino heating and eventually leads to neutrino-driven explosion. This is a new scenario of core-collapse supernovae that has never been demonstrated by numerical simulations so far. ",
        "introduction": "At present neutrino heating is considered to be the most promising candidate for the as-yet-unknown explosion mechanism of core-collapse supernovae (CCSNe). Recent state-of-the-art simulations with detailed neutrino transport have not yet succeeded in reproducing the canonical explosion energy of $10^{51}$~erg, though \\citep[e.g.,][]{suw10,mue12,bru13}. Other effects to enhance or drive explosion may hence be required.  Magnetic field and rotation may bring such effects. Although the magnetic field and rotation of progenitors at the pre-collapse stage are not well-understood, some of recent theoretical and observational studies indicate that their strengths may spread over a wide range.  Optical observations of OB main sequence stars reported recently that some of them possess $\\sim 1$~kG surface magnetic fields \\citep[e.g.,][]{wad12}, which correspond to the magnetar-class magnetic flux, $\\sim 10^{26}$--$10^{27}$~G cm$^{2}$. A population synthesis calculation assuming the conservation of magnetic flux during the post-main-sequence evolution performed by \\citet{fer06} indicates that $\\sim 10$\\% of OB stars have the magnetar-class magnetic flux, while the majority has $\\sim$1--2 orders of magnitude weaker ones.  Stellar evolution calculations so far indicate that the pre-collapse rotation rate is sensitive to that at the zero-age main sequence (ZAMS). Whereas \\citet{heg05} found that a 15~$M_\\odot$ star with the solar metallicity and the surface rotational velocity of $v_{\\textrm{rot,ZAMS}}=$200~km~s$^{-1}$ at ZAMS results in the pre-collapse rotation that may be translated to a pulsar rotation period of $P_{\\textrm{NS}}=11$~ms, a 16~$M_\\odot$ star with the solar metallicity and $v_{\\textrm{rot,ZAMS}}=$360~km~s$^{-1}$ computed by \\citet{woo06} ends up with a pre-collapse rotation corresponding to $P_{\\textrm{NS}}=2.3$--9.7~ms, depending on the unknown mass-loss rate in the Wolf-Rayet stage.  Meanwhile, \\citet{ram13} found that 20~\\% of 216 O-type stars in 30 Dor have the surface rotational velocities larger than 300~km~s$^{-1}$. Combined with the results of \\citet{woo06}, this implies that rapidly rotating progenitors that may produce a pulsar rotation period of milliseconds, may not be so rare even for the solar metallicity.  The influences of magnetic field and rotation on the dynamics of CCSNe have been so far studied mostly in the regime in which magnetic field is large and rotation is rapid simultaneously, i.e., the combination of the magnetar-class magnetic flux and the rotation that would produce millisecond proto-neutron star (MPNS) has been assumed. Numerical simulations have shown that the magnetic field, which is already quite large initially and is later amplified by differential rotation, can drive explosions with the canonical energy of $10^{51}$~erg \\citep[e.g.,][]{yam04,obe06,shi06,saw13a}.  There are a small number of simulations assuming initially sub-magnetar-class magnetic flux. In their axisymmetric (2D) simulations \\citet{bur07} and \\citet{tak09} found that even with sub-magnetar-class fields, $B_{\\textrm{pre}}\\lesssim 10^{11}$~G at pre-collapse, rapid rotations corresponding to MPNS can amplify the magnetic fields to produce well-collimated magneto-driven jets. On the other hand, similar simulations by \\citet{moi06} obtained more or less spherical explosions. In both cases dipole-like magnetic fields assumed at pre-collapse generate explosion energy smaller than the canonical value. Note, however, that all these simulations do not have enough spacial resolutions to capture magnetorotational instability (MRI), which is expected to occur in CCSNe and, if true, would drastically amplify the magnetic field in the timescale of rotation \\citep{bal91,aki03}.  Resolving MRI in CCSNe with sub-magnetar-class magnetic flux demands quite a fine size of numerical grid with a width of, say, a few 10 m for $B_{\\textrm{pre}}\\sim 10^{11}$~G, which should be compared with the size of iron cores, $\\sim 1000$~km. To deal with high computational cost, some previous simulations were done in local boxes placed in PNSs \\citep{obe09,mas12}. \\citet{saw13b}, on the other hand, performed global simulations in axisymmetry, deploying as many as $9300\\times 6400$ mesh points and demonstrated that MRI produces strong magnetic fields on large scales even under the dynamical background. They also showed that the magnetic force becomes dynamically important in some locations. Since their simulations were terminated around 70~ms after bounce, the consequence for the dynamics thereafter were not clear. Note that neutrino transport was not included in their simulations, which is main reason they stopped computations.  In this letter, we investigate the aftermath. In so doing, we assume again the initial magnetic fields with sub-magnetar-class magnetic flux and carry out long-term MHD simulations in axisymmetry with neutrino cooling and heating, albeit simplified, being taken into account.  In the following, the numerical methods and models are described in \\S\\ref{sec.model}, results are presented in \\S\\ref{sec.result}, and discussion and conclusion are given in \\S\\ref{sec.conc}. ",
        "conclusions": "\\label{sec.conc} We conducted two dimensional high-resolution global MHD simulations of CCSNe, assuming that a massive star core is rather weakly magnetized and rapidly rotating prior to collapse. Taking neutrino heating and cooling into account approximately, we followed the evolutions long after core bounce and studied the dynamical consequences of MRI. We found that the magnetic field amplified by MRI indeed affects the dynamics thereafter substantially enhances the explosion. In fact, it efficiently transfers angular momentum from high to low latitudes and expands the heating region at low latitudes not by magnetic stress but by centrifugal force. This then enhances the heating efficiency and leads to explosion. This is a new explosion mechanism, which we expect to work generally in weakly magnetized, rapidly rotating stellar cores. In fact, although even our highest-resolution run has not yet achieved numerical convergence particularly in the sub-dominant poloidal component, further improvement of resolution would result in more efficient angular momentum transfer and thus in higher neutrino heating, which would make shock revival by neutrino heating even easier.  Previous simulations of weakly magnetized, rapidly-rotating cores by \\citet{bur07} and \\citet{tak09} found magneto-driven jets formed along the rotation axis, which are analogous to what we found in the L-MRI run. Since the spacial resolution in \\citet{bur07} is in between our BG and L-MRI runs, where the transition from mass accretion to ejection occurs near the pole, their results do not contradict ours. What they observed may be an artifact by insufficient numerical resolutions as demonstrated in the current study. Although \\citet{bur07} claimed that the neutrino heating is subdominant (contributing 10--25~\\% to the explosion energy), our simulations indicate that it is actually predominant. The spacial resolution in \\citet{tak09} is coarser than our BG run. The reason for the jet formations in their simulations may be because they assumed strong differential rotation prior to collapse with a typical scale height of 100~km in the $\\varpi$ direction.  The diagnostic explosion energy obtained in our highest resolution simulation (H-MRI run) is $\\sim 10^{49}$~erg, which is far smaller than the canonical value, $10^{51}$~erg. It should be noted that the shock front reaches only the radius of 500~km at the end of the simulation. It is way too early to estimate the final value of the explosion energy, since it may increase as the shock front propagates outward. Further improvement of the spacial resolution may also increase the explosion energy, since more efficient angular momentum transfer is expected. Moreover, the neutrino luminosity of $10^{52}$erg s$^{-1}$ assumed in this study may be a bit too low. In fact, a recent simulation by \\citet{bru13} obtained $\\nu_e$/$ \\bar{\\nu_e}$ of several $10^{52}$erg s$^{-1}$ for the first 200~ms after bounce, which decreases to $\\approx 10^{52}$erg s$^{-1}$ only later at $\\sim 700$~ms after bounce. Larger neutrino luminosities will certainly result in larger explosion energies. Apart from the quantitative estimation of the explosion energy, what is most important here is the finding that even a weak magnetic field added to a rapidly rotating core may make the neutrino-driven explosion easier."
    },
    "1402/1402.5931_arXiv.txt": {
        "abstract": "{Cygnus A, as the nearest powerful FR II radio galaxy, plays an important role in understanding jets and their impact on the surrounding intracluster medium.} {To explain why the nucleus is observed superposed onto the eastern lobe rather than in between the two lobes, and why the jet and counterjet are non-colinear. } {We made a comparative study of the radio images at different frequencies of Cygnus A, in combination with the published results on the radial velocities in the Cygnus A cluster.} {From the morphology of the inner lobes we conclude that the lobes are not interacting with one another, but are well separated, even at low radio frequencies. We explain the location of the nucleus as the result of the proper motion of the galaxy through the cluster. The required proper motion is of the same order of magnitude as the radial velocity offset of Cygnus~A with the sub-cluster it belongs to. The proper motion of the galaxy through the cluster likely also explains the non-co-linearity of the jet and counterjet.} {}  \\keywords {galaxies:active--galaxies:individual: Cygnus~A--galaxies:jets}  \\titlerunning{} \\authorrunning{K. C. Steenbrugge et al.} ",
        "introduction": "}  Cygnus~A (3C~405) has been well studied at radio wavelengths as it was one of the first double sources detected at radio wavelengths (see for instance \\citealt{carilli91} and \\citealt{carilli96}). At GHz frequencies and above a jet directed at $\\sim$60$^{\\circ}$ to our line of sight, embedded within the western lobe, is observed. A weaker counterjet approximately anti-parallel to the jet is observed within the counterlobe or eastern lobe. In a set of previous papers we have studied the jets and lobes of Cygnus~A using radio and X-ray images and X-ray spectra. We showed that the jets are precessing, have a best fit speed of 0.35~c, and that the angle between the jet and inner counterjet is 179$^{\\circ}$, determined from the best-fit precession model \\citep{steenbrugge08a}. From a visual inspection of the jet and counterjet, these authors determined a somewhat smaller, angle of 177$^{\\circ}$$\\pm$1$^{\\circ}$.5. This result is unexpected and remains unexplained. 40~kpc out from the nucleus the counterjet bends through a 27$^{\\circ}$ 31$^{\\prime}$ angle, the cause of which is still poorly understood, but is possibly related to the presence of relic plasma \\citep{steenbrugge10}.  We also found that the plasma from a previous episode of jet activity explains the presence of a relic counterjet in the X-ray image of Cygnus~A \\citep{steenbrugge08b} and the fact that at low frequencies the counterlobe is {\\it brighter} than the lobe \\citep{steenbrugge10}. Although a well known general result, important for this article is that the size of the lobe and counterlobe do depend on the observed frequency. Both the lobe and counterlobe are much larger at 151~MHz than at 5~GHz and are larger at 5~GHz than at 15~GHz (see for instance figure~5 in \\citealt{steenbrugge10}).  This radio galaxy is embedded in the Cygnus cluster formed by the merging of 2 clusters, each probably having a richness class 1 \\citep{markevitch99,ledlow05}. The ongoing merger is clearly observed in the X-ray temperature map and supported by the radial velocity clustering and location of the galaxies \\citep{markevitch99,belsole07,ledlow05}. From the radial velocity structure, the location of the galaxies, and the observed X-ray temperature structure, \\cite{ledlow05} estimate that core encounter will occur in 0.2$-$0.6~Gyr hence. The average radial velocity of the cluster galaxies identified by \\cite{ledlow05} is 19,008$\\pm$160~km~s$^{-1}$ with a velocity dispersion of 1126~km~s$^{-1}$. For Cygnus A these authors measured a radial velocity of 16,811~km~s$^{-1}$, and thus an offset of 2197~km~s$^{-1}$ from the cluster average. However, the measured radial velocities are clustered in two groups, one with an average radial velocity of 17,648~km~s$^{-1}$ and one with an average radial velocity of 20,126~km~s$^{-1}$ (as determined by \\citealt{ledlow05} using location and radial velocity information). From its radial velocity, the host of this powerful FR~II classical double radio galaxy is at a redshift of 0.05607 \\citep{owen97}. Assuming a Hubble constant of 73 km s$^{-1}$ Mpc$^{-1}$, 1$^{\\prime\\prime}$ corresponds to 1.044~kpc.  Cygnus~A is not in the centre of either of the subclusters \\citep{owen97}, and actually shows an offset of 72 $h^{-1}_{75}$~kpc from the centre of the smaller subcluster \\citep{ledlow05}. The centres of the two subclusters are separated by 458~$h^{-1}_{75}$~kpc \\citep{ledlow05}. There is an overall temperature difference discerned in the cluster, due to the merger. At the location of Cygnus~A, but mostly on larger scales than the radio lobes, an asymmetry in temperature is clearly observed \\citep{belsole07}. Due to the difference in timescale between the cluster merger and the jet age of the current activity, and the fact that the asymmetry in the intracluster medium (ICM) is on larger scales, the radio lobes are not affected by the ongoing merger process.  In this paper we study the morphology of the region surrounding the nucleus using radio images at different frequencies. In aligning the radio and X-ray data, it is clear that the nucleus of the Cygnus~A galaxy does not lie between both lobes, but actually lies in front of, within or behind the counterlobe. Explaining this result, we find a likely explanation for a second observational result, namely that the jet and inner counterjet are not co-linear \\citep{steenbrugge08a}.  In the following Section we summarize the data used for this analysis. In Section~3 we describe the morphology of the inner lobes and the nucleus. The discussion is given in Section~4 and our conclusions are stated in Section~5. ",
        "conclusions": "In this article we studied the position of the nucleus with respect to the lobes as well as the trajectory of the inner jets. We found that the nucleus is located in front of or in the counterlobe, the best explanation for which is a proper motion of the Cygnus~A galaxy through the Cygnus cluster, probably caused by the ongoing merger in this cluster. The derived range in proper motion, between 60 and 900~km~s$^{-1}$ has the same order of magnitude as the range in radial velocity offsets determined by \\cite{ledlow05}. The lack of bending over larger angles as seen in WAT radio sources seems to indicate that Cygnus~A is located in a rather low density region compared to central cluster densities. This proper motion is also likely to be the explanation that the observed angle between the inner jet and counterjet with the nucleus is 179$^{\\circ}$ as determined from the precession modelling. Future, deep and high-resolution low-frequency polarisation images should be able to determine whether the nucleus is located in or in front of the counterlobe."
    },
    "1402/1402.0725_arXiv.txt": {
        "abstract": "{We present evidence that spots imaged using astronomical CCDs do not exactly scale with flux: bright spots tend to be broader than faint ones, using the same illumination pattern. We measure that the linear size of spots or stars, of typical size 3 to 4 pixels FWHM, increase linearly with their flux by up to 2 \\% over the full CCD dynamic range. This brighter-fatter effect affects both deep-depleted and thinned CCD sensors. We propose that this effect is a direct consequence of the distortions of the drift electric field sourced by charges accumulated within the CCD during the exposure and experienced by forthcoming light-induced charges in the same exposure. The pixel boundaries then become slightly dynamical: overfilled pixels become increasingly smaller than their neighbors, so that bright star sizes, measured in number of pixels, appear larger than those of faint stars. This interpretation of the brighter-fatter effect implies that pixels in flat-fields should exhibit statistical correlations, sourced by Poisson fluctuations, that we indeed directly detect. We propose to use the measured correlations in flat-fields to derive how pixel boundaries shift under the influence of a given charge pattern, which allows us in turn to predict how star shapes evolve with flux. We show that, within the precision of our tests, we are able to quantitatively relate the correlations of flat-field pixels and the broadening of stars with flux. This physical model of the brighter-fatter effect also explains the commonly observed phenomenon that the spatial variance of CCD flat-fields increases less rapidly than their average.} ",
        "introduction": "\\label{sec:intro} The astronomical use of CCDs very commonly assumes that these sensors slice the incoming light along a regular lattice of pixels, each of which then responds independently of what adjacent pixels receive. Non-linearity usually refers to the reponse of the electronic chain, which can be (and often is) corrected independently for each pixel.  In this representation, the images scale precisely with the incoming flux or staring time on a static scene, possibly after some linearity restoration is applied at the pixel level. We will discuss in this contribution a tiny violation of this idealistic picture of CCDs: images of stars (or laboratory spots) do not exactly scale with the incoming flux, even with a prefectly linear readout chain.  With increasing flux, images of spots or stars tend to broaden: the measured sizes grow linearly with flux. Since the effect is almost isotropic (though not exactly), there is a strong inclination to attribute it to sensors rather than to the read out chain. All attempts to detect the effect we know about have been successful, and the effect does not seem to only affect deep-depleted CCDs.  There is another known departure from the independent pixels representation: the variance of a flat-field image does not grow linearly with the exposure average (as Poisson statistics robustly predict), but the variance ``flattens out'', as if some mechanism washing the contrast were at play. This phenomenon is sufficiently similar to the brighter-fatter effect to attempt to bridge these two effects.  The flattening of the photon transfer curve (PTC i.e. the spatial variance of flat-fields as a function of their average) is known to be reduced or to almost disappear by summing neighboring pixels into larger pixels (e.g. \\cite{Downing06}). This indicates that the variance of a sum of neighboring pixels is larger than the sum of their variances, and adjacent pixels should therefore have (positive) covariances. We detect these covariances. They contain more information than the flattening of the PTC, and we will hence study pixel correlations in flat-field images rather than the ``missing variance'' one can extract from the shape of the PTC.  Large scale imaging surveys, underway and anticipated, vitally require precise knowledge of their Point Spread Function (PSF). Regarding the accuracy of the PSF shape and size, the most demanding application is arguably the measurement of the shear of background galaxy images induced by mass density gradients along the line of sight, in order to constrain these mass gradients, and in turn cosmology. Usually, the design of the data reduction chain assumes that stars collected in the science image are faithful models of the PSF, so that the effects of the observing system (possibly including atmosphere) on objects shapes can be measured and ``factored out'' using these stars. We will argue that stars are acceptable PSF templates only for applications that can accommodate an inaccuracy of the PSF size at the percent scale. In particular, studying the correlations of the shear field measured from galaxy shapes now requires a PSF typically understood at the $10^{-3}$ level or better (see e.g. the contribution of M. Jarvis at this workshop\\footnote{The presentations at this workshop can be found on \\href{http://www.bnl.gov/cosmo2013/}{http://www.bnl.gov/cosmo2013/}}). The requirements on PSF ellipticity accuracy are as low as $2\\ 10^{-4}$ for the Euclid space mission \\cite{EuclidRB}.  We attempt in this contribution to bridge the brighter-fatter effect and the correlations in flat-fields, assuming that both are a consequence of alterations of drift lines in the CCD bulk induced by charges accumulating in the potential wells of the sensor during the exposure. These drift field alterations can be regarded as dynamical modifications of pixel boundaries\\footnote{The ``tree-rings'' phenomenon (see the contributions of e.g. G. Bernstein, E. Magnier \\& R. Lupton at this workshop) is probably due to a similar phenomenon, but induced by a static distortion of the drift field sourced by spatial and static variations of the silicon bulk resistivity.}.  We wish to set up this bridge in order to derive the details of the brighter-fatter effect from pixel correlations in flat-fields. This would allow one to account for the brighter-fatter effect in the data reduction procedures, for a variety of image qualities and PSF shapes, and beyond the growth of second moments.   In this paper we proceed as follows. First we expose the characteristics of the brighter-fatter effect (\\S~\\ref{sec:detection}). We then discuss statistical pixel correlations in flat-fields (\\S~\\ref{sec:correlations}). We propose in \\S~\\ref{sec:calculations} that Coulomb forces sourced by already collected charges on drifting ones are the source of both effects, and propose a parametrized model that explains both. In \\S~\\ref{sec:from-flat-to-bf}, we derive the model parameters from flat-field images and then predict the brighter-fatter slope for three different sensor types. We summarize and discuss possible further work in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} Coulomb forces within CCDs are a plausible cause of the brighter-fatter effect and statistical correlations observed in flat-fields. The latter are also related to the flattening of the PTC commonly observed with CCD devices. We propose a scheme to describe the brighter-fatter effect at the pixel level using measured correlations in flat-fields that matches measurements at a typically $\\sim$20\\% accuracy.  Further work is required before we can practically handle the brighter-fatter effect along this proposed line. We have neglected the fact that charges in the CCD not only alter drift lines but also reduce the longitudinal drift field (this is usually referred to as ``space charge effect'').  As a result, the measured correlations in flat-fields are slightly increased by the effect, generally absent in science data because the sky level is most often deliberately kept low by adjusting the exposure time. This decrease of the longitudinal drift field also slightly increases the lateral diffusion for bright spots with respect to faint ones (see the contribution by S. Holland at this workshop). However, the size of this effect is much smaller than what we observe, in particular for the practical case of a PSF broader than 3 pixels FWHM. Among the three sensor types we have discussed, the one from DECam is the most affected by these space-charge effects because it has the lowest nominal drift field: the bias voltage is 40 V for a thickness of 250~$\\mu$m.  Properly handling the brighter-fatter effect aims at evaluating the PSF shape at all fluxes where it is practically needed. The non-linearity of response of the sensor (typically due to the electronic chain) can also produce variations of the PSF size with flux, and has to be handled separately. In our comparisons, we have ignored possible non-linearities. For Megacam, there is practical evidence that they are very small (see Appendix C in \\cite{Regnault09}). For the E2V CCD data, variations of the light source intensity with time limit our capability to address linearity. For DECam, a significant non-linearity has been detected (see the contribution of G. Bernstein at this workshop) by regressing the average measured level of dome flat-fields against exposure time. One might wonder if the electronic chain responds in a similar way to high steady fluxes (as for high-level flat-fields) and to localized high signal levels in a few pixels (as for bright stars). Taking these non-linearity measurements at face value increases the observed non-linearity but the disagreement between measured and simulated bright-fatter slopes remains of similar size, with the opposite sign.  One obvious weakness of the proposed approach is that the general model has more coefficients than the number of two-point flat-field correlations we can measure. One possible way out is to improve electrostatic simulations in order to derive relations between model coefficients which are reasonably independent of the unknowns affecting the sensor layout. Regarding measurements, we have not settled yet if higher order flat-field correlations deliver new information compared with two-point correlations, nor if higher order correlations are practically measurable. One can also consider measuring the response of sensors to other patterns than just flat-fields to constrain the evolution of shapes of sharp spots. How the brighter-fatter effect varies with image quality likely contains some amount of information.  Regarding the practical handling of the brighter-fatter effect in reduction pipelines, we can already imagine two avenues: the simple one consists in ``de-scrambling'' the images at the pixel level, using some inverse of relation \\ref{eq:full-model}. If just flipping signs of the $a$ coefficients is not accurate enough, some refinement can easily be devised. However, this model assumes that the image is properly sampled, i.e. that the flux on pixel boundaries is well approximated by interpolating between the pixel contents. This is likely to be an unacceptable approximation for sharp images as typically obtained from space. One can then account for the brighter-fatter effect inside the PSF model, using the PSF itself to evaluate the charge density along pixel boundaries.  It is likely that the crude technique we have exposed in this work can handle the brighter-fatter effect at the $\\sim$10\\% accuracy, with some refinements, possibly along the lines mentioned in the above paragraphs. It is certainly premature to consider that a 1\\% accuracy is within reach. In particular, the brighter-fatter effect is anisotropic, because the pixel boundaries separating line and rows of CCDs are generated via totally different mechanisms for most CCD types. We believe that the way to account for the evolution of PSF ellipticity with flux at the $\\sim 10^{-4}$ level, as required for very large scale weak shear programs (see \\cite{EuclidRB}), is still an open question."
    },
    "1402/1402.2666_arXiv.txt": {
        "abstract": "We present \\nustar hard X-ray % observations of three \\typeii quasars at $z\\approx0.4$--$0.5$, optically selected from the Sloan Digital Sky Survey (SDSS). Although the quasars show evidence for being heavily obscured Compton-thick systems on the basis of the $2$--$10$~keV to \\oiii luminosity ratio and multiwavelength diagnostics, their X-ray absorbing column densities ($N_{\\rm H}$) are poorly known. In this analysis: (1) we study X-ray emission at $>10$~keV, where X-rays from the central black hole are relatively unabsorbed, in order to better constrain $N_{\\rm H}$; (2) we further characterize the physical properties of the sources through broad-band near-UV to mid-IR spectral energy distribution (SED) analyses. One of the quasars is detected with \\nustar at $>8$~keV with a no-source probability of $<0.1\\%$, and its X-ray band ratio suggests near Compton-thick absorption with $N_{\\rm H} \\gtrsim 5 \\times 10^{23}$~\\nhunit. The other two quasars are undetected, and have low X-ray to mid-IR luminosity ratios in both the low energy ($2$--$10$~keV) and high energy ($10$--$40$~keV) X-ray regimes that are consistent with extreme, Compton-thick absorption ($N_{\\rm H}\\gtrsim 10^{24}$~\\nhunit). We find that for quasars at $z\\sim 0.5$, \\nustar provides a significant improvement compared to lower energy ($<10$~keV) \\chandra and \\xmm observations alone, as higher column densities can now be directly constrained. ",
        "introduction": "\\label{Introduction}  Quasars are the sites of the most rapid black hole growth in the universe \\citep{Salpeter64,Soltan82}. They represent the luminous end of the active galactic nucleus (AGN) population, often outshining their host galaxies. The first unobscured (`\\typei') quasars were discovered over 50 years ago \\citep{Schmidt63, Hazard63}, and more than one hundred thousand have now been spectroscopically identified (e.g., \\citealt{VeronCetty10, Paris12}).  For obscured (`\\typeii') quasars\\footnote{We define \\typeii quasars as AGN with $L_{\\rm 2-10keV} \\geq 10^{44}$~\\ergpersec, X-ray absorbing column densities $N_{\\rm H}>10^{22}$~\\nhunit, and optical spectra that show narrow line emission without broad (H$\\rm \\alpha$ or H$\\rm \\beta$) components.  This $L_{\\rm 2-10keV}$ threshold is consistent with: (1) the classical optical quasar definition, $M_{B}\\leq -23$, when the $\\alpha _{\\rm OX}$ relation of \\citet{Steffen06} and the composite quasar spectrum of \\citet{VandenBerk01} are assumed; (2) the $L_{\\rm X,*}$ value derived by \\citet{Hasinger05} for unobscured AGN.} the situation is not as advanced.  Similar to the early \\typei quasars, \\typeii quasars were initially identified from radio selection (e.g., \\citealt{Minkowski60}), and over the following decades several hundred powerful `radio galaxies' (as such radio-selected \\typeii\\ quasars are typically called) were identified (for reviews, see \\citealt{McCarthy93, Miley08}).  However, it is only in the past decade that radio-quiet \\typeii quasars have been found in large numbers.  Such sources are generally identified on the basis of either their relatively hard X-ray spectral slopes (e.g., \\citealt{Norman02, Stern02}), optical spectral features (e.g., \\citealt{Steidel02, Zakamska03}), or mid-infrared (mid-IR) colors (e.g., \\citealt{Lacy04, Stern05}).  Importantly, mid-IR color selection of \\typeii quasars using the all-sky {\\it Wide-Field Infrared Survey Explorer} ({\\it WISE}; \\citealt{Wright10}) survey identifies several million \\typeii quasars, roughly down to the bolometric luminosity of the primary Sloan Digital Sky Survey \\citep[SDSS;][]{York00} \\typei quasar spectroscopic survey \\citep{Stern12, Assef13, Donoso13}.  The exact nature of \\typeii quasars is still under debate. A simple extension of the orientation-driven unified model of AGN \\citep{Antonucci93, Urry95} to high luminosities can account for their existence. However, there is also observational evidence for an evolutionary link to \\typei quasars \\citep[e.g.,][]{Sanders88, Hopkins08}.  The importance of \\typeii quasars to the cosmic evolution of AGN is further demonstrated by their requirement in models of the cosmic X-ray background (CXB) \\citep[e.g.,][]{Treister05, Gilli07, Treister09}.  However, the observed X-ray properties of \\typeii quasars are poorly constrained at present. Consequently, the column density ($N_{\\rm H}$) distribution\\footnote{X-rays emitted from the immediate black hole environment are absorbed by circumnuclear gas, and thus provide constraints on $N_{\\rm H}$.} and Compton-thick\\footnote{Compton-thick absorption is that with $N_{\\rm H} \\ge {\\rm \\sigma}_{\\rm T}^{-1}\\approx1.5\\times 10^{24}$~\\nhunit.} fraction of quasars are poorly known, which has implications for both AGN and CXB models \\citep[e.g.,][]{Fabian08,Draper10}.    To date, the largest sample of spectroscopically confirmed (radio-quiet) \\typeii quasars at $z\\lesssim1$ is that of \\citet{Zakamska03} and \\citet{Reyes08}. \\citet{Zakamska03} selected 291 \\typeii quasars at redshift $0.2 \\lesssim z \\lesssim 0.8$ from the SDSS based on their optical properties: high \\oiii $\\lambda 5007$ line power and narrow emission lines. \\citet{Reyes08} used the same approach and more recent SDSS data to extend the sample to 887 objects. While X-ray selections of \\typeii quasars at $\\lesssim 10$~keV are biased against the most heavily obscured sources \\citep[e.g.,][]{Maiolino98a}, \\oiii emission is mostly produced on $\\sim100$~pc scales and is thus relatively unaffected by nuclear obscuration, allowing larger numbers of the heavily obscured, X-ray faint objects to be found. Following up \\oiii selected, rather than X-ray selected, objects with X-ray observations thus gives a less biased estimate of the $N_{\\rm H}$ distribution of AGN \\citep[e.g.,][]{Risaliti99}.    The X-ray properties of the \\citet{Zakamska03} and \\citet{Reyes08} \\typeii quasar sample have been studied using \\chandra and \\xmm observations \\citep{Ptak06,Vignali06,Vignali10,Jia13}. \\citet{Vignali06,Vignali10} measured column densities for a handful of sources through `direct' means (i.e., using X-ray spectroscopic analysis). The highest column densities measured in this manner were $N_{\\rm H}\\approx 3 \\times 10^{23}$~\\nhunit. However, distant obscured quasars are X-ray weak and in most cases direct constraints are not feasible. Instead, an `indirect' approach to estimating column densities can be used where the observed X-ray emission is compared with a proxy for intrinsic AGN power \\citep[e.g., the mid-IR continuum emission from hot dust or high-excitation emission lines;][]{Bassani99,Lutz04,Heckman05,Alexander05b,Alexander08,Cappi06,Panessa06,Melendez08,Gandhi09,LaMassa09,LaMassa11,Gilli10,Goulding11}. \\citet{Vignali06,Vignali10} were limited to indirect absorption constraints for the majority of their \\typeii quasar sample, and found in every case that Compton-thick absorption (i.e., $N_{\\rm H} > 1.5 \\times 10^{24}$~\\nhunit) is required to explain the X-ray suppression in these sources. To first order, there appears to be a bimodal $N_{\\rm H}$ distribution for optically selected \\typeii quasars, with $\\sim 40\\%$ having $N_{\\rm H} = 10^{22}$--$3\\times10^{23}$~\\nhunit and $\\sim 60\\%$ being Compton-thick. This is interesting given that a continuous $N_{\\rm H}$ distribution is measured for % \\typeii Seyferts \\citep[e.g.,][]{Bassani99,Risaliti99,LaMassa09,LaMassa11}, although the differences may be reconciled by considering the different methods used to estimate $N_{\\rm H}$ \\citep{LaMassa11}. To better constrain the $N_{\\rm H}$ distribution of \\typeii quasars, more robust identifications of Compton-thick absorption must be obtained through either: (i) measurement of strong \\feka emission, with ${\\rm EW}\\geq1$~keV, which results from the \\feka line being viewed in reflection against a suppressed continuum \\citep[e.g.,][]{Ghisellini94,Levenson02}; % or (ii) measurement of high column densities through spectroscopic analysis at high energies above the photoelectric absorption cutoff (i.e., above observed-frame $8$~keV for $z\\sim 0.5$ and $N_{\\rm H}\\sim10^{24}$~cm$^{-2}$), where X-ray emission is relatively unabsorbed.     The recent launch of the {\\it Nuclear Spectroscopic Telescope Array} \\citep[\\nustar,][]{Harrison13} will see a breakthrough in our understanding of heavily obscured AGN and the CXB population in general. \\nustar is the first orbiting observatory with the ability to focus high energy ($\\gtrsim 10$~keV) X-rays using grazing incidence optics. It provides a two orders of magnitude improvement in sensitivity and over an order of magnitude improvement in angular resolution over previous hard X-ray observatories. The high energy range at which \\nustar operates ($3$--$79$~keV) means that the intrinsic, unabsorbed emission of AGN is observed for all but the most heavily obscured, Compton-thick objects. At $z \\lesssim 1$, it is now possible to directly constrain column densities an order of magnitude higher than those achievable with \\chandra and \\xmm alone \\citep[e.g.,][]{Luo13}.  In this paper, we present exploratory \\nustar observations of three optically selected \\typeii quasars at $z \\approx 0.4$--$0.5$. All three have been identified as Compton-thick candidates in previous studies \\citep{Vignali06,Vignali10,Jia13}. We use X-ray data from \\nustar, \\chandra and \\xmm, and near-UV to mid-IR data from other observatories to determine the physical properties of the quasars. In particular, we use a combination of direct and indirect methods to constrain the absorbing column densities. The paper is organized as follows: our sample selection is detailed in Section \\ref{sample}; we describe the observations, data reduction and data analysis in Section \\ref{data}; % our main results regarding X-ray absorption constraints are presented in Section \\ref{Results}; we summarize our main conclusions in Section \\ref{summary}. The cosmology adopted throughout this work is $(\\Omega_{M},\\Omega_{\\Lambda},h)=(0.27,0.73,0.71)$. ",
        "conclusions": ""
    },
    "1402/1402.6790_arXiv.txt": {
        "abstract": "{We report the detection, with {\\it Chandra} and XMM-{\\it Newton}, of faint, soft X-ray emission from four symbiotics stars that were not known to be X-ray sources. These four object show a $\\beta$-type X-ray spectrum, i.e. their spectra can be modeled with an absorbed optically thin thermal emission with temperatures of a few million degrees. Photometric series obtained with the Optical Monitor on board XMM-{\\it Newton} from V2416 Sgr and {NSV 25735} support the proposed scenario where the X-ray emission is produced in a shock-heated region inside the symbiotic nebulae.} ",
        "introduction": "\\label{sec:intro}  In a binary system where a white dwarf (WD) accretes from the wind of a red giant companion and frequently forms an accretion disk; where the WD can experiences frequent outburst of different intensity and where mass is ejected at speeds of a few hundreds km s$^{-1}$, it seems natural to expect that X-ray emission could be frequently detected, originated either in the surface of the white dwarf by quasi-stable nuclear burning, in the accretion disk, or in shocks due to ejections of material. However, it was only recently that a significant fraction of these systems, known as WD symbiotics, was detected at X-ray wavelengths using modern instruments such as {\\it Swift}, XMM-{\\it Newton} or {\\it Chandra}. In general, the emission is faint, with fluxes as low as a few 10$^{-14}$ ergs cm$^{-2}$ s$^{-1}$. WD symbiotics have X-ray spectra that can be classified into four types: supersoft X-ray sources whose X-ray spectrum peaks at energies lower than 0.4 keV, named $\\alpha$-type; $\\beta$-type comprises those sources with soft X-ray emission and whose spectrum extends up to energies of less than 2.4 keV;  $\\delta$-type are highly absorbed, hard X-ray sources with detectable thermal emission above 2.4 keV and $\\beta/\\delta$-type includes those WD symbiotics with two X-ray thermal components, soft and hard \\citepads[this classification scheme is detailed in][hereafter Paper I]{2013A&A...559A...6L}.  In this paper we present the detection at X-ray energies of four WD symbiotics: \\object{Hen~2-87}, \\object{NSV 25735}, \\object{V2416~Sgr} and \\object{Hen~2-104} observed with XMM-{\\it Newton} and {\\it Chandra} that were not known to be X-ray sources before. In Section \\ref{sec:analisis} we detailed the data reduction while in Section \\ref{sec:results} we present the results, while discussion and concluding remarks are presented in Section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We found X-ray emission from four WD symbiotic that were not known to be X-ray sources. The X-ray emission from all of them can be classified as $\\beta$-type, following the scheme originally proposed by \\citetads{1997A&A...319..201M} and recently updated in Paper I. Our results confirm that $\\beta$-type X-ray emission is the most frequent among WD symbiotic, increasing the number of known $\\beta$-type sources to 16 out of 48 symbiotic with X-ray emission. The X-ray emission from the new sources is faint, with luminosities on average of 10$^{31}$ ergs s$^{-1}$ (albeit the uncertainties in their distance), which explains why pointing, dedicated observations were necessary for their detection.  As $\\beta$-type systems, their X-ray emission seems to be associated with shocks in the nebulae that could be originated in a colliding wind region or inside extended jets (known to be present in \\object{NSV 25735} and \\object{Hen~2-104}).  Although most models of colliding-wind scenarios require that the WD in the system underwent a recent outburst to be able to trigger a fast, tenuous wind, the $\\beta$-type systems presented here as well as those presented in Paper I do not have any recent outburst recorded. While $\\alpha$ or $\\delta$ type X-ray spectra are directly linked to the source that powers the WD, i.e. shell-burning or accretion, $\\beta$ type X-ray emission is not. The absence of strong variability, together with a $\\beta$-type X-ray spectrum strongly suggest that the observed X-rays are not originated in the accretion disk around the WD, while at the same time, those soft X-rays are too energetic to be originated in a shell-burning white dwarf. Our physical picture is that, in these systems, nuclear burning could lead to a strong wind from the WD (or the accretion disk near the WD), which collides with the wind from the cool star, leading to observed soft X-ray emission. Super-soft X-ray emission from the WD might be present as well but absorbed by the dense symbiotic nebula.  Using the temperature obtained from spectral models of these sources (0.4 keV to $\\sim$2.0 keV) and assuming strong shock conditions, we can derive a shock speed, v$_{shock}$. We obtain v$_{shock}$ in the range of $\\sim$1\\,000 to $\\sim$2\\,500 km s$^{-1}$. Similar speeds are observed in the $\\beta$-type sources presented in Paper I and \\citetads{1997A&A...319..201M} and are roughly consistent with the speeds of X-ray emitting outflows from WD symbiotic \\citepads{2007ApJ...660..651N,2004ApJ...613L..61G}. High outflow speeds of a few thousands km s$^{-1}$ have been observed in the optical line profiles of MWC~560 \\citepads{2001A&A...377..206S}, although they are not consistent with the temperatures observed from the $\\beta$ component in its X-ray spectrum \\citepads{2009A&A...498..209S}. The disparity between X-ray infered speed and speeds from optical line widths has been also observed in the diffuse X-ray emission from planetary nebulae \\citepads{2003ApJ...583..368S}. In the dense symbiotic nebula, however, some authors have modeled the observed UV line profiles of O~VI and He~II as due to electron scattering \\citepads{2012BaltA..21..196S,2012MNRAS.427..979S}. \\citetads{2000ApJ...541L..25L} proposed that Ramman scattering might be responsible for the broad H$\\alpha$ line wings observed in symbiotics.  The number of known symbiotic with X-ray emission is rapidly increasing with the new instruments available and it is expected to increase even more with much sensitive instruments and deeper surveys (e.g. {\\it NuSTAR}\\footnote{http://www.nustar.caltech.edu/} and {\\it eROSITA}\\footnote{http://www.mpe.mpg.de/eROSITA}). We now know 48 systems with X-ray emission originating in four different scenarios. With this number of sources it seems timely to study their statistical properties. Is there any trend or relationship between IR and X-ray spectral types?. In the symbiotic stars catalog of \\citetads{2000A&AS..146..407B}, 63 \\% of the objects (139) are classified as $S$-type, 15\\% (34) as $D$-type, and 3\\% (7) as $D'$-type. Among the 48 symbiotic with X-ray emission, 33\\% of them have a $\\beta$-type X-ray spectrum. Roughly half of the $\\beta$-type symbiotics have infrared $S$-type spectrum, while 25\\% have $D$-type, 6.3\\% have $D'$-type and 6.3\\% do not have classification. We did not find any clear tendency or correlation between IR and X-ray types, i.e. the fractions of IR types are roughly similar among X-rays emitting symbiotics and those with no X-ray emission detected."
    },
    "1402/1402.1453_arXiv.txt": {
        "abstract": "We study the impact of the existence of an eV-mass scale sterile neutrino---with parameters in the ballpark of what is required to fit the laboratory anomalies---on the early time profile of the electron neutrino and antineutrino fluxes associated to a core-collapse supernova (SN). In particular, we focus on the universal feature of neutronization burst expected in the first tens of ms of the signal: provided that a detector with sufficient sensitivity is available, it is well-known that in the 3 neutrino framework the detection of the neutronization burst in neutrino channel would signal inverted mass hierarchy. This conclusion is dramatically altered in the presence of a sterile neutrino: we study here both analytically and numerically the region in parameter space where this characteristic signal disappears, mimicking normal hierarchy expectations. Conversely, the detection of a peak consistent with expectations for inverted mass hierarchy would exclude the existence of a sterile state over a much wider parameter space than what required by laboratory anomalies fits, or even probed by detectors coming on-line in the near future. Additionally, we show the peculiar alteration in the energy-time double differential flux, with a delayed peak appearing for kinematical reasons, which might offer a remarkable signature in case of favorable parameters and for a high statistics detection of a Galactic SN. We also comment on additional potentially interesting effects in the electron antineutrino channel, if more than one angle in the active-sterile sector is non-vanishing. As an ancillary result that we derived in the technical resolution of the equations, in an appendix we report the Cayley-Hamilton formalism for the evolution of a four neutrino system in matter, generalizing existing results in the literature. ",
        "introduction": "A very exciting frontier of low-energy neutrino ($\\nu$) astronomy is represented by the detection of neutrinos from core-collapse supernovae (SNe), whose first and till now only example was provided by SN1987A~\\cite{Hirata:1987hu,Hirata:1988ad,Bionta:1987qt}. Existing large underground neutrino detectors (like SuperKamiokande or IceCube) as well as numerous planned ones are well suited to detect these rare galactic events (a few per century {\\it in average}) with sufficiently high-statistics to allow for the extraction of detailed astrophysical information on the SN explosion mechanism~\\cite{GilBotella:2004bv,Pagliaroli:2008ur,Pagliaroli:2009qy,Ianni:2009bd,Halzen:2009sm,Lund:2010kh}. Such measurements could also offer a handle on particle physics such as $\\nu$ masses and mixings, too~\\cite{Dighe:1999bi,Lunardini:2001pb,Lunardini:2003eh,Dighe:2003be,Abbasi:2011ss,Serpico:2011ir}.  One of the open questions in the neutrino sector is the existence of light sterile neutrino states, motivated by some experimental ``anomalies\". Although a three (active) neutrino mixing scenario explains most of the data consistently, some short-baseline neutrino oscillation experiments suggest some deviations. These include the ${\\overline \\nu}_\\mu \\to {\\overline\\nu}_e$ oscillations in LSND~\\cite{Aguilar:2001ty} and MiniBooNE~\\cite{Aguilar-Arevalo:2013pmq} experiments, the ${\\overline\\nu}_e$ and $\\nu_e$ disappearance dubbed Reactor Anomaly~\\cite{Mention:2011rk} in reactor neutrino experiments and the Gallium Anomaly~\\cite{Acero:2007su,Giunti:2010zu} in the calibration of solar neutrino experiments, respectively. Interestingly, a large fraction of these outlier data can be roughly accommodated if a light ($m\\sim\\mathcal O(1)\\,$eV) sterile neutrino state is added to the picture (the so-called $3+1$ model)~\\cite{Kopp:2013vaa,Giunti:2013aea,Giunti:2012tn}, although tensions between different sets of data persist. A plethora of experiments, relying on different strategies and methods, are in design or construction phase to check the existence of sterile neutrinos (see~\\cite{Abazajian:2012ys} and references therein; see also~\\cite{Esmaili:2013vza,Esmaili:2013cja,Esmaili:2012nz,Esmaili:2012vg}).  Not surprisingly, the existence of such a sterile state would have an impact on SN neutrino conversions. Recently, for example, Ref.~\\cite{Tamborra:2011is} showed how in such conditions the electron fraction $Y_e$ is reduced and the conditions for heavy-element formation in the supernova ejecta can be affected. It is also clear that the signal detectable at the Earth can be altered, see for instance~\\cite{Choubey:2006aq,Choubey:2007ga} for an example associated to the turbulent shock wave.  In this article we discuss another interesting effect, which has passed almost unnoticed till now. In particular, the existence of sterile states with properties in the ball-park of what is required by the interpretation of ``anomalies\" lead to interesting phenomenological consequences on the {\\it neutronization burst}, that should be observable in a large detector of $\\nu_e$. The neutronization burst is a prompt burst of $\\nu_e$ associated to the passage of the newly formed shock to regions with densities low enough that neutrinos (initially trapped) begin to stream faster than the shock. Since the medium is basically made of free nucleons and is rich in $e^-$, the only rapid process is $e^-$ captures on $p$: this significantly suppresses the flux of flavors other than $\\nu_e$, while the $\\nu_e$ signal lasts. The existence and time profile of this burst are a generic feature of the early signal (first ${\\cal O}$(20) ms post-bounce) of core-collapse SNe. Its properties are largely independent of the progenitor and still uncertain physical properties such as the dense matter equation of state (see for example~\\cite{Kachelriess:2004ds} or Fig.~1 in~\\cite{Serpico:2011ir}.) The detection of such a burst has already been discussed in the literature as a way to establish inverted mass hierarchy (IH) in the (active) neutrino sector~\\cite{GilBotella:2003sz,Kachelriess:2004ds}, possibly  the most robust and spectacular one from SN $\\nu$, provided that an instrument with enough sensitivity is available. In fact, for the presently measured ``large'' value of $\\theta_{13}$, in the neutrino channel maximal $\\nu_e$ conversion occurs for normal hierarchy (NH), while the survival probability is constant and given by $P_{ee}\\simeq \\sin^2 \\theta_{12}\\cos^2{\\theta_{13}}\\simeq 0.32$ for IH (see e.g.~\\cite{GilBotella:2003sz} for details). Note that detectors such as a ${\\cal O}(100)\\,$kton Liquid Argon time projection chamber---of the same class proposed within the context of the LAGUNA collaboration for a future underground detector~\\cite{Autiero:2007zj}---or a Megaton class water Cherenkov detector~\\cite{Kachelriess:2004ds}, have already been shown to be capable of such a measurement of neutronization burst signal.  This ``unambiguous'' picture is significantly altered in presence of a sterile neutrino with parameters fitting the laboratory anomalies. Notably, to first approximation the existence of sterile neutrino can make the $\\nu_e$ burst signature disappear altogether! This fact has been mentioned in the past, see e.g.~\\cite{Peres:2000ic,Murayama:2000hm}, but has never been studied in detail. Here we present a more precise analytical and numerical discussion of this signature: In particular, we identify the region in parameter space where the phenomenon takes place, comparing them to the preferred ones from sterile neutrino explanations of laboratory anomalies. Additionally, we present the peculiar alteration in the energy-time double differential flux, highlighting the presence of a delayed ``peak'' appearing for {\\it kinematical} reasons, which may be a non-negligible feature for a sufficiently large mixing angle and mass in the sterile sector. The phenomenological importance of these features for diagnostics both in the active and sterile sector is discussed. We also study the consequences of assuming more than one non-vanishing angle in the active-sterile sector: this is particularly important for potential signatures in the $\\bar{\\nu}_e$ detection channel, which we briefly address.   This article is structured as follows: The formalism of $3+1$ scenario and discussions about the resonant active-sterile neutrino conversion is reported in Sec.~\\ref{formalism}. In Sec.~\\ref{FluxComp} we report the SN $\\nu_e$ flux composition in $3+1$ scenario and in Sec.~\\ref{sec:flux} we illustrate the effect of the existence of a sterile neutrino state with typical mixing parameters on representative time-energy double differential SN $\\nu$ fluxes. Sec.~\\ref{phenonuebar} is devoted to the SN $\\overline{\\nu}_e$ flux, its composition at Earth and phenomenological considerations. In Sec.~\\ref{discussion} we shall discuss our results and finally conclude with some perspectives for forthcoming studies.  As a side remark of some technical importance, it is worth pointing out that we checked our analytical results (which assume factorization and adiabaticity) with a numerical code which implements the evolution of a $4\\nu$ system with appropriate parameters in a (toy) SN matter potential, which is the generalization of the method described in~\\cite{Ohlsson:1999xb} based on the Cayley-Hamilton formalism. Since we could not find the explicit result in the literature, we worked out the relevant formulae and report them in Appendix~\\ref{4x4CH}. For completeness, in Appendix~\\ref{FluxEarth} we also report the details of the derivation of our analytical results for the SN $\\nu_e$ and $\\overline{\\nu}_e$ flux compositions outside the SN surface as a function of the input fluxes at the neutrinosphere. ",
        "conclusions": "\\label{discussion}   The next Galactic supernova explosion and the observation of its neutrino flux in the existing and forthcoming experiments at Earth will provide a unique opportunity to study both the explosion mechanism and neutrino physics. In this paper we discussed how the existence of a fourth, mostly sterile neutrino state $\\nu_4$ (heavier than the active ones, the so-called $3+1$ model) with a $\\nu_e-\\nu_s$ mixing characterized by the mixing element $U_{e4}$ would alter the expected $\\nu_e$ SN flux at Earth. Obviously, the effect depends on the mass of new state $m_4$ and its mixing $U_{e4}$. However, for a wide range of parameter values (see Eqs.~(\\ref{resupp}), (\\ref{eq:NHrange}), and (\\ref{eq:IHrange})) the  $\\nu_e$ radiated from the neutrinosphere convert completely to the $\\nu_4$ state en route to the surface of the SN, and hence to the detector at the Earth. Since $\\nu_4$ is mostly sterile, this resonant conversion drastically alters the expected early time neutronization $\\nu_e$ burst, making it unobservable at leading order. In more detail, due to the small (but not necessarily negligible) mixing $|U_{e4}|^2$, the $\\nu_4$ flux has still a chance to be detectable as $\\nu_e$ on Earth. However, its kinematic characteristics are altered: depending on the mass $m_4$, energy $E_\\nu$ and the distance of SN to Earth, $D$, the $|U_{e4}|^2$-proportional $\\nu_e$ flux will be delayed by a time $D(m_4/E_\\nu)^2/2c$. We provided an analytical description of the relevant physics, and checked our analytical results (which assume $2\\times2$ factorization and adiabaticity of the resonances) against numerical calculations, finding a good agreement. The numerical computations were performed with a $4\\times 4$ generalization of the Cayley-Hamilton formalism described in~\\cite{Ohlsson:1999xb}, and we report the relevant formulae in Appendix~\\ref{4x4CH}.   Our main results can be thus summarized as follows: If the mass hierarchy will be unknown at the time of future Galactic SN detection, the presence of a fourth sterile state can fake the NH phenomenology (lack of observable neutronization burst) even for IH in the active neutrino sector. Turning the argument around, should the active neutrino hierarchy be determined to be of the IH type, the existence of a sterile state may be one of the simplest explanation for  a lack of visible neutronization burst observation from a future SN events in a sufficiently large $\\nu_e$ detector. This may corroborate independent evidence from the lab, but also be sensitive to mixing values below current constraints. On the other hand (and perhaps more important), the observation of a neutronization peak consistent with expectations for IH would exclude the existence of a sterile state over a much wider parameter space than what required by laboratory anomalies fits, or even the one testable by detectors coming on-line in the near future. This provides yet another nice example of interplay and complementarity of the astroparticle observables with laboratory ones.   What are the chances that this signature can be actually observed? In the past decade, there have been dedicated studies concerning the detectability of the neutronization burst with different techniques, see notably~\\cite{GilBotella:2003sz,Kachelriess:2004ds}. Here we just recall the main results, requirements and challenges, addressing the reader to the original literature for details.  Obviously, the identification of the neutronization burst is especially clean with detectors using the charged-current absorption of $\\nu_e$'s. The most widely discussed large detector option for this channel is provided by liquid argon, dominantly via $\\nu_e+{}^{40}{\\rm Ar}\\to e^{-}\\,+{}^{40}{\\rm K}^*$. The study in~\\cite{GilBotella:2003sz}, considering a 70 kton detector, showed that the presence or absence of a neutralization burst leads to a count number of events as different as 86 vs. 41 within the first 240 ms of the signal, for a fiducial SN model located at 10 kpc from us (note that from within 10 kpc  one expects roughly 50\\% chances to observe the next SN, see for example the distribution in Ref.~\\cite{Mirizzi:2006xx}). Even accounting for Poisson fluctuations and (small) model-to-model variations, in Ref.~\\cite{Kachelriess:2004ds} it was estimated that a $2\\sigma$ discrimination could be achieved already by this counting test, provided that the distance to the SN is known. Needless to say, a closer SN could allow a separation even with a smaller detector (or equivalently to a higher confidence level, for the benchmark case of 70 kton), while an uncertainty in the position would worsen the sensitivity. Note however that a more refined test exploiting the {\\it time structure} might improve the perspectives for diagnostics.  The other technique that has been investigated concerns large water Cherenkov detectors, such as the proposed Hyper-Kamiokande in Japan. This experimental technique is less clean, since the $\\nu_e$ elastic scattering on electrons has to compete with other large signals from inverse beta decay on protons, reactions on oxygen and $\\nu_x$ scattering onto electrons, but the larger masses (Mton scale) permit to take advantage of the higher statistics. Furthermore, the $\\nu_e$ elastic scattering on electrons are more forward peaked and less energetic than most background events. The other channels can be most effectively separated if the detector is doped with Gadolinium, as suggested in~\\cite{Beacom:2003nk} and currently tested at the EGADS facility with encouraging results~\\cite{EGADS}. Accounting for statistical errors, nuclear cross section uncertainties, supernova model dependence (such as progenitor mass, equation of state) it was found that the capability of distinguishing the two cases are often better than $3\\sigma$ for a fiducial SN at 10 kpc, and never worse than $2\\sigma$.  In conclusion, testing for the presence of a conventional neutrino burst in the next Galactic SN signal appears within the reach of next generation of underground neutrino detectors, and actually providing a further particle physics motivation to tackle these major experimental enterprises. It is worth noting, however, that it may be possible to obtain a detection of the neutronization peak already with the currently operating IceCube detector at the South Pole. This instrument offers ``only'' a calorimetric light curve via the correlated increase of ``Cherenkov noise\" in its detectors (and thus typically via the inverse beta decay reaction), but the statistics is so high that fine time structures can be revealed. It has been shown that in the first 30 ms or so post-bounce ({\\it i.e.}, of emergence of a signal on top of the instrumental noise) the two scenarios with/without neutralization peak are markedly different, see e.g. Fig.~11 in~\\cite{Abbasi:2011ss}.  The observational perspectives for the other signature (delayed and energy-distorted peak) remain to be studied. Its detection requires at least comparable performances as for the neutronization peak detection, if not superior, plus some luck in the particle physics parameters, such as relatively large mixing angles and large masses. Energy and timing resolution also play a great role. One can envisage in fact to optimize specific strategies to exploit the peculiar time-energy correlation, perhaps extending earlier proposals for the ``active neutrino'' mass measurements from SN signals, see e.g.~\\cite{Totani:1998nf}. For sure, the detection of the delayed peak would be a very specific signature of this scenario, but a dedicated analysis is needed to explore the observational perspectives in the allowed parameter space. We note here that other exotic phenomena have been discussed in relation to the prompt neutronization burst: for example, the appearance of the burst in the $\\bar{\\nu}_e$ channel due to magnetic moments~\\cite{Akhmedov:2003fu} or neutrino decay~\\cite{Ando:2004qe}. If anything, these possibilities should highlight the importance of large underground detectors for a high-statistics measurement of the neutrino flux(es) from the next Galactic core-collapse SN.  While being probably the most spectacular one, the alteration of the neutronization burst is not the only manifestation of the presence of sterile neutrinos in the expected neutrino fluxes. We briefly discussed how the antineutrino channel would also be altered, in particular if small mixings of the sterile state with the $\\nu_\\mu$ or $\\nu_\\tau$ are present. Most neutrino detectors use the inverse beta decay reaction for detection and thus are sensitive primarily to $\\bar{\\nu}_e$, so that this channel may offer a more easily accessible diagnostic tool. We showed how large (up to a factor 3!) alterations of the appearance probabilities are induced by the presence of a sterile state. A natural follow-up of our article would be to study the observational signatures of this channel as well, either in the number, energy and time distribution of events in a Water Cherenkov detector, say, or in the luminosity profile that can be measured with impressive detail in a detector like IceCube. Finally, one might wonder if specific signatures of the kinematical time-delay may be inferred from other techniques than the study of the neutronization burst. One possible direction would be to consider if alterations of the time variation of the neutrino emissions revealed in simulations (at ms level, due to anisotropic mass flows in the accretion layer around the newly-formed neutron star) are detectable, along the lines of the study~\\cite{Ellis:2012ji} for mass constraints of the active neutrinos."
    },
    "1402/1402.1179_arXiv.txt": {
        "abstract": "Imaging Atmosphere Cherenkov Telescopes (IACT) are arrays of very large optical telescopes that are well-suited for rapid photometry of bright sources.  I investigate their potential in observing stellar occultations by small objects in the outer Solar System, Transjovian Objects (TJOs).  These occultations cast diffraction patterns on the Earth.  Current IACT arrays are capable of detecting objects smaller than 100 {\\Blue metres} in radius in the Kuiper Belt and 1 km radius out to 5000 AU.  The future Cherenkov Telescope Array (CTA) will have even greater capabilities.  Because the arrays include several telescopes, they can potentially measure the speeds of TJOs without degeneracies, and the sizes of the TJOs and background stars.  I estimate the achievable precision using a Fisher matrix analysis.  With CTA, the precisions of these parameter estimations will be as good as a few percent.  I consider how often {\\Blue detectable} occultations {\\Blue occur} by members of different TJO populations, including Centaurs, Kuiper Belt Objects (KBOs), Oort cloud objects, and satellites and Trojans of Uranus and Neptune.  The great sensitivity of IACT arrays means that they likely detect KBO occultations once every ${\\cal O}(10)$ hours when looking near the ecliptic.  IACTs can also set useful limits on many other TJO populations. ",
        "introduction": "\\label{sec:Introduction}  There are many minor bodies in the Solar System beyond the orbit of Jupiter, the Transjovian Objects (TJOs).  There are several populations of TJOs.  These include the Centaurs, a collection of objects that orbits between Jupiter and Neptune; the Kuiper Belt, a reservoir of bodies such as Pluto orbiting 30 - 50 AU from the Sun; the Scattered Disc Objects, a lower density population (among them, Eris; {\\Blue \\citealt*{Brown05}; \\citealt{Brown07}}) that extends out to 100 - 200 AU {\\Blue \\citep*[e.g.,][]{Gladman08}}; and the Oort Cloud, a population of comets that is generally believed to mostly reside tens of thousands of AUs from the inner Solar System, but includes objects like Sedna that orbit hundreds of AU away {\\Blue \\citep*{Brown04}}.  The TJOs formed out of the debris left over after planet formation, and their physical properties potentially contain information about the early Solar System.  Their orbits also record the effects of gravitational perturbations from the planets, and constrain models of Solar System dynamics \\citep*[e.g.,][]{Morbidelli08}.  Because the TJOs are in the distant reaches of the Solar System, our ability to observe and understand them is limited.  Large TJOs can be directly observed with telescopes; following the early discovery of Pluto, other large TJOs were found with surveys starting in the 1990s (e.g., \\citealt{Jewitt93}; {\\Blue \\citealt{Brown04,Brown05,Kavelaars08,Trujillo08}}).  These big objects have been studied intensely with photometry and spectroscopy {\\Blue \\citep[e.g.,][]{Lazzarin03,Peixinho04,Stransberry08}}.  Pluto itself, {\\Blue and possibly} more distant Kuiper Belt objects, will be visited by the \\emph{New Horizons} probe, which will increase our understanding of TJOs {\\Blue \\citep{Stern08}}.\\footnote{Triton, the largest moon of Neptune, is likely a captured Kuiper Belt object \\citep[e.g.,][]{Agnor06}, and was visited by \\emph{Voyager 2} in 1989.  However, its structure was radically altered by its capture \\citep*[e.g.,][]{McKinnon95}.}  However, TJOs with radii of a kilometre or less are far more difficult to observe.  General constraints on the amount of mass in small bodies can be found from the level of CMB anisotropies (\\citealt*{Babich07}; \\citealt{Ichikawa11}), the brightness of the infrared background \\citep{Kenyon01}, and the gamma-ray background \\citep{Moskalenko08,Moskalenko09}.  The main method of searching for these {\\Blue small} objects is through stellar occultations, when they pass in front of a background star, blocking and diffracting its light (\\citealt{Bailey76,Dyson92}; \\citealt*{Roques87}; \\citealt{Roques00}).  {\\Blue At least two TJOs were detected with this method in the past few years \\citep{Schlichting09,Schlichting12}.}  The Fresnel scale for objects {\\Blue in} the outer Solar System is $\\ell_F = \\sqrt{\\lambda D / 2} = 1.4\\ \\km\\ (\\lambda / 5500\\ \\Ang)^{1/2} (D / 50\\ \\AU)^{1/2}$.  Objects larger than $\\ell_F$ essentially cast a geometrical shadow on the Earth, while smaller objects cast a diffraction pattern on the Earth as wide as $\\ell_F$, with the star's magnitude fluctuations set by the size of the occulter.  Observing these events poses several challenges.  First, the depth of the fluctuations for the smaller TJOs is typically only of order a few percent.  Second, because the Earth moves with a relative speed of up to $\\sim 30\\ \\kms$ with respect to TJOs, the events last only a fraction of a second.  High quality light curves must be therefore sampled with a frequency 5 -- 100 Hz {\\Blue (cf., \\citealt{Nihei07}, hereafter N07)}.  Past and present surveys for TJO occultations include TAOS, the Taiwanese American Occultation Survey, which is dedicated to detecting TJO occultations \\citep{Lehner09}; additionally, surveys have been carried out on the MMT observatory \\citep{Bianco09}, the Panoramic Survey Telescope and Rapid Response System 1 (PS1) telescope \\citep{Wang10}, the Very Large Telescope \\citep{Doressoundiram13}, and Hubble Space Telescope \\citep{Schlichting09}.  Additional searches for occultations of Sco X-1 in the X-ray band have been {\\Blue conducted}, but instrumental effects proved troublesome (\\citealt{Chang06,Chang07,Jones08,Liu08}; \\citealt*{Chang11}).  Historically, a big problem with the occultation method in visible light is scintillation noise, which dominates fluctuations at short frequencies.  The technique therefore requires large telescopes with high frequency sampling, and must deal with scintillation by either having a large signal-to-noise ratio \\citep{Wang10,Doressoundiram13}, using multiple telescopes as a veto \\citep{Bianco10,Lehner10}, or {\\Blue using} space telescopes to eliminate it altogether \\citep{Schlichting09}.  Many of these conditions are fulfilled by the Imaging Atmospheric Cherenkov Telescopes (IACTs), which are among the largest optical telescopes in the world.  Cherenkov telescopes are used to detect TeV gamma-rays by imaging the flash of Cherenkov light emitted by particle showers created when the gamma ray hits the upper atmosphere (e.g., {\\Blue \\citealt{Galbraith53};} \\citealt{Aharonian97}).  The flashes are faint, so the Cherenkov telescopes must be large to collect as many photons as possible (see Table~\\ref{table:IACTArrays}). In addition, the Cherenkov flashes are very short, only a few nanoseconds long, so Cherenkov telescopes use photomultipliers or silicon photon detectors to sample fluxes on MHz time-scales.  The optical field of view of a Cherenkov telescopes is a few degrees. Finally, Cherenkov telescopes often come in arrays, {\\Blue allowing} them to image a particle shower in three dimensions, but {\\Blue this} could also be useful for vetoing occultation false positives.  Cherenkov telescopes achieve this remarkable performance at a low cost by sacrificing angular resolution: the typical {\\Blue point spread function (PSF)} of a Cherenkov telescope is a few arcminutes \\citep{Bernlohr03,Cornils03}.  This increases the noise of a source because of blending with more sky background.  However, for the brightest sources (with $V \\la 8 - 10$), the Poisson noise of the star is greater than the sky background noise, and confusion only becomes a problem at $V \\approx 14 - 16$ \\citep{Bahcall80}.  Therefore, the Cherenkov telescopes may be very useful {\\Blue for observing} bright stars.  Cherenkov telescopes have been used to study the Crab pulsar optical light curve \\citep{Hinton06,Lucarelli08}, optical transients on millisecond and microsecond time-scales \\citep{Deil09}, and for optical SETI \\citep{Eichler01,Holder05}, and they are potentially good for stellar intensity interferometry \\citep{LeBohec06}, detecting picosecond optical transients \\citep{Borra10} and optical polarimetry \\citep{Lacki11}.  Current Cherenkov telescope arrays include the Very Energetic Radiation Imaging Telescope Array System (VERITAS)\\footnote{http://veritas.sao.arizona.edu/} and the High Energy Stereoscopic System (HESS)\\footnote{http://www.mpi-hd.mpg.de/hfm/HESS/} (see Table~\\ref{table:IACTArrays}).  VERITAS consists of four 12-meter telescopes and is located in the Northern hemisphere \\citep{Weekes02}.  HESS is an array of four 12-meter telescopes plus a central, fifth 28-meter telescope, located in the Southern hemisphere \\citep{Bernlohr03,Cornils03,Cornils05}.  The power of Cherenkov telescopes will increase much further with the next generation Cherenkov Telescope Array (CTA).\\footnote{http://www.cta-observatory.org/}  The array will include four large ($\\sim 25\\ \\meter$) telescopes, as well as tens of medium ($\\sim 12\\ \\meter$) and small ($\\sim 7\\ \\meter$) telescopes \\citep{Bernlohr13}.  The sheer number of telescopes will allow a vast number of photons to be collected, and would efficiently eliminate false positives.  The precise observations of stellar occultations that are possible with Cherenkov telescopes can be useful for studying not only the TJOs themselves but the occulted stars.  The shape of the light curve depends on the size of the TJO, as well as the angular size of the star, in units of the Fresnel scale.  An array of telescopes, by sampling the light curve at different positions, {\\Blue could} measure the size of the diffraction pattern, giving the Fresnel scale, from which these parameters and the TJO distance can be calculated.  Besides these parameters, the arrays can measure the position of the TJO on the sky, and the non-radial components of the TJO velocity.  \\begin{table*} \\begin{minipage}{110mm} \\caption{TJO occultation instrument configurations} \\label{table:IACTArrays} \\begin{tabular}{lccccc} \\hline Array   & $N_{\\rm tel}$ & $D_{\\rm tel}$ Range & $A_{\\rm tel}$ & Baseline Range & References\\\\ &               & $m$                 & $m^2$         & $m$            & \\\\ \\hline TAOS         & 4        & 0.5                 & 0.79          & 7.6 -- 100     & (1)\\\\ Megacam/MMT  & 1        & 6.5                 & 33            & ...            & (2)\\\\ ULTRACAM/VLT & 1        & 8.2                 & 53            & ...            & (3)\\\\ \\hline VERITAS      & 4        & 12.0                & 450           & 83 -- 180      & (4)\\\\ HESS         & 5        & 12.0 -- 28.0        & 1100          & 85 -- 170      & (5)\\\\ CTA (E)      & 59       & 7.2 -- 24.0         & 5800          & 67 -- 3500     & (6)\\\\ \\hline \\end{tabular} \\\\{\\bf References} -- (1) \\citet{Lehner09}; (2) -- as used in \\citet{Bianco09}; (3) -- as used in \\citet{Doressoundiram13}; (4) -- \\citet{Weekes02}, present configuration given in \\citet*{Perkins09}; (5) -- see \\citet{Bernlohr03}, \\citet{Cornils05}, and the HESS website, http://www.mpi-hd.mpg.de/hfm/HESS; (6) -- configuration E given in \\citet{Bernlohr13} \\\\ \\end{minipage} \\end{table*}  In \\citet{Lacki11}, I computed the expected signal-to-noise ratios for photometric, spectroscopic, and polarimetric measurements by Cherenkov telescopes given the large sky noise.  Noting that these ratios were better for bright sources, I listed several possible phenomena that might be studied optically with Cherenkov telescopes, including stellar occultations by TJOs.  Now I evaluate the abilities of Cherenkov telescopes to detect and characterize these occultations.  After describing how I calculate the light curves of occultation events in Section~\\ref{sec:Calculations}, I evaluate the types of events that are detectable at VERITAS, HESS, and CTA in Section~\\ref{sec:Detectability}.  The ability of Cherenkov telescopes to estimate the parameters of an occultation is evaluated with Fisher matrix analysis {\\Blue and likelihood ratio analysis} in Section~\\ref{sec:ParameterEstimation}.  Finally, the frequency that these arrays should observe TJO occultations is calculated in Section~\\ref{sec:Frequency}. ",
        "conclusions": "IACT arrays are in many ways superior instruments for studying TJO occultations.  Their telescopes are enormous with great light collection areas and relatively little scintillation noise; the telescopes come in arrays, allowing a veto of scintillation noise and potentially breaking the velocity-size-distance degeneracy; the photon sampling rates are essentially infinite from the occultation perspective; and their fields of view ($\\sim 5^{\\circ}$) are large enough to observe dozens of stars at once.  The disadvantages are large amounts of sky noise because of the poor angular resolution, too short baselines for TJO parameter estimation, and relatively low duty cycle because of moonlight.  {\\Blue In addition, the current electronics of the PMTs in IACTs may not be suited for continuous photometry; if they do not use the standard readout modes, there can be deadtime within the integration time (Benbow, private communication).}  I have shown that the already extant VERITAS and HESS arrays are in principle capable of detecting occultations by sub-kilometre TJOs.  The two KBO occultations detected by the HST would have been easily detected by either of these arrays.  HESS, with its central 28 {\\Blue metre} telescope, can detect more objects than VERITAS, but has comparable performance for parameter estimation.  {\\Blue The greatest obstacle for these arrays is estimating $\\ell_F$.}  I estimate that they observe a KBO occultation of a $V = 12$ star once every $\\la 10$ hours of near-ecliptic observations. The arrays {\\Blue could} also detect a number Centaurs and Scattered Disc Objects during the course of a year, and are capable of setting interesting constraints on the number of Neptune Trojans and irregular satellites.  Much rarer occultations of blue $V = 10$ stars occur at $\\la 10\\%$ of the rate of the $V = 12$ stars, but observations of these events are worthwhile, as they allow for precise determinations of the TJO size and distance.  The future CTA will be an excellent facility for characterizing KBO occultations as it will include kilometre-long baselines.  It should be able to detect objects with radii smaller than 100 {\\Blue metres} in the Kuiper Belt, allowing it to explore the same parameter space as X-ray occultation searches \\citep[e.g.,][]{Liu08,Chang11}, but with many more targets.  CTA also can detect 1 km radius objects out to several thousand AUs, into the inner reaches of the Oort Cloud.  Even for an occultation of a $V = 12$ star by a 250 {\\Blue metre} radius KBO, CTA will be able to constrain properties to within a few percent.  Given the apparent great potential of these instruments, it would be worth considering more complex models.  First, a better model of scintillation noise is desirable, as its properties for large telescopes are poorly understood, and it may be correlated on relevant time-scales \\citep{Bickerton09}.  Second, the IACTs are so powerful that the assumptions of a uniformly bright star and a spherical object may be insufficient.  A study of whether these instruments can constrain limb darkening, starspots, and the shape of the star and the occulting body should be carried out.  Finally, my estimates of the occultation rates are very crude; a more advanced model of the stellar population as in \\citet{Cooray03-Rates} would be useful."
    },
    "1402/1402.6335_arXiv.txt": {
        "abstract": "We report the detection at 850\\,$\\mu$m of the central source in SSA22-LAB1, the archetypal `Lyman-$\\alpha$ Blob' (LAB), a 100\\,kpc-scale radio-quiet emission-line nebula at $z=3.1$. The flux density of the source, $S_{\\rm 850}=4.6\\pm1.1$\\,mJy implies the presence of a galaxy, or group of galaxies, with a total luminosity of $L_{\\rm IR}\\approx10^{12}\\,{\\rm L_\\odot}$. The position of an active source at the center of a $\\sim$50\\,kpc-radius ring of linearly polarized Ly$\\alpha$ emission detected by Hayes et al.\\ (2011) suggests that the central source is leaking Ly$\\alpha$ photons preferentially in the plane of the sky, which undergo scattering in H{\\sc i} clouds at large galactocentric radius. The Ly$\\alpha$ morphology around the submillimeter detection is reminiscent of biconical outflow, and the average Ly$\\alpha$ line profiles of the two `lobes' are dominated by a red peak, expected for a resonant line emerging from a medium with a bulk velocity gradient that is outflowing relative to the line center. Taken together, these observations provide compelling evidence that the central active galaxy (or galaxies) is responsible for a large fraction of the extended Ly$\\alpha$ emission and morphology. Less clear is the history of the cold gas in the circumgalactic medium being traced by Ly$\\alpha$: is it mainly pristine material accreting into the halo that has not yet been processed through an interstellar medium (ISM), now being blown back as it encounters an outflow, or does it mainly comprise gas that has been swept-up within the ISM and expelled from the galaxy? ",
        "introduction": "Lyman-$\\alpha$ Blobs (LABs) are $\\sim$100\\,kpc scale radio-quiet Lyman-$\\alpha$ emission line nebulae, with integrated luminosities up to $L_{\\rm Ly\\alpha}\\sim10^{37}$\\,W. LABs were first discovered during the course of deep optical narrowband imaging of the Small Selected Area 22$^{\\rm hr}$ (SSA22) field (Steidel et al.\\ 2000), which was known to contain a significant overdensity of Lyman-break Galaxies (LBGs) at $z=3.09$ (Steidel et al.\\ 1998), although previously Francis et al.\\ (1996) and Keel et al.\\ (1999) had also detected extended Ly$\\alpha$ emission around luminous galaxies at $z=2.4$ (both in high density environments) and extended Ly$\\alpha$ emission was reported to be associated with some of the first high-{\\it z} submillimetre galaxies (SMGs, e.g., Ivison et al.\\ 1998). Subsequent narrowband (and broadband, see Prescott et al.\\ 2012, 2013) surveys of the SSA22 `protocluster' and elsewhere have since revealed a substantial sample of these high redshift nebulae, with a range of Ly$\\alpha$ luminosities and extents (e.g.,\\ Matsuda et al.\\ 2004, 2011, Yang et al.\\ 2009, Erb et al.\\ 2011, Bridge et al.\\ 2013) and often inhabiting dense environments (e.g.\\ Prescott et al.\\ 2008).  Since their discovery, the nature of LABs has been in debate. It was recognised early on that the galaxy halo cooling budget would contain a significant cold component (White \\& Frenk\\ 1991), and low metallicity gas at $T\\sim10^4$--$10^5$\\,K cools most efficiently via Ly$\\alpha$ emission through collisional excitation (Fabian \\& Nulsen\\ 1977, Katz \\& Gunn\\ 1991, Katz et al.\\ 1996, Haiman et al.\\ 2000), so one train of thought is that extended Ly$\\alpha$ emission implies the presence of potentially pristine cold gas in the circumgalactic medium (CGM). Taking this idea further, many cosmological hydrodynamical simulations suggest that the baryonic growth of massive ($M_{\\rm h}>10^{12}\\,{\\rm M_\\odot}$) galaxies at $z>2$ (and for lower mass galaxies at {\\it all} epochs) is dominated by a mode of `cold accretion' in the form of narrow filamentary streams that penetrate the hot, virially shocked halo gas (Katz et al.\\ 2003, Kere\\v{s} et al.\\ 2005,\\ 2009, Dekel et al.\\ 2009). Nevertheless, the predicted cold flows have so-far eluded unambiguous observation and some recent simulations seem to refute the existence of cold flows that survive to the disc altogether (Nelson et al.\\ 2013). Nevertheless, LABs have been touted as the closest we have come to identifying the filamentary cold mode accretion phenomenon in nature.  While the cosmological gas simulations can track the history of cold gas easily enough, the predicted fluxes of the cooling Ly$\\alpha$ emission are far more uncertain, with the flux depending on the complex interplay of turbulence, radiative transfer and local sources of ionizing radiation are taken into account (e.g.,\\ Rosdahl \\& Blaziot\\ 2012). The prescriptions for Ly$\\alpha$ radiative transfer through the gas are notoriously complex due to the fact that it is a resonant line (Neufeld\\ 1990) and sensitive to assumptions about the sub-grid physics governing the transport of Ly$\\alpha$ photons through astrophysical media, as well as the many other possible sources of Ly$\\alpha$ photons in and around galaxies. As a result, different simulations that apply bespoke models of Ly$\\alpha$ radiative transfer, aiming to reproduce the properties of observed LABs, predict a wide range of physical properties (Ly$\\alpha$ luminosities can vary by an order of magnitude) and key observables, such as Ly$\\alpha$ extents and surface brightnesses (Faucher-Gigu\\`ere et al.\\ 2010, Goerdt et al.\\ 2010, Rosdahl \\& Blaziot 2012).   The cold flow picture is complicated further when one considers the effects of the presence of a luminous source (a starburst galaxy or one hosting an active galactic nucleus, AGN) embedded within the LAB, which appears to be a common situation (e.g.,\\ Francis et al.\\ 1996, Keel et al.\\ 1999, Dey et al.\\ 2005, Geach et al.\\ 2005,\\ 2007,\\ 2009, Colbert et al.\\ 2006, Webb 2009, Rauch et al.\\ 2011, Cantalupo, Lilly \\& Haenhelt\\ 2012). The typical bolometric luminosities of the embedded sources imply an energy balance heavily weighted towards the galaxy, such that only relatively weak coupling, through feedback, would be required to explain the Ly$\\alpha$ luminosities (Geach et al.\\ 2009). Again, there is further uncertainty in the type of feedback that would give rise to such extended Ly$\\alpha$ emission: proposals have included emergent (supernovae-driven) superwinds that shock heat cold gas swept up in the outflow (Taniguchi \\& Shioya\\ 2000), photoionization of the CGM by escaping ultraviolet continuum radiation from massive stars or the accretion discs of an AGN, or photoionization via inverse Compton (IC) scattering of cosmic microwave background (or locally generated far-infrared) photons, which are energetically up-scattered by a relic population of relativistic electrons (e.g.,\\ Scharf et al.\\ 2003, Fabian et al.\\ 2009).  Geach et al.\\ (2009) rule-out IC scattering as a viable LAB formation mechanism using a deep X-ray stack in SSA22, arguing that photoionization by a central AGN (which may be present in up to 50\\% of LABs) is probably more important, even with fairly modest escape fractions likely to be applicable to the often heavily obscured sources. Recently, Hayes et al.\\ (2011) presented evidence that at least some of the extended Ly$\\alpha$ emission must be in the form of scattered Ly$\\alpha$ radiation, with Ly$\\alpha$ photons generated within a central galaxy scattering in clouds of H{\\sc i} at large ($\\sim$50\\,kpc) galactocentric radius. The absence of an obvious luminous counterpart in some LABs (e.g.,\\ Nilsson et al.\\ 2006, Smith \\& Jarvis\\ 2007, Smith et al.\\ 2008) has been inferred to be a sign that, in some cases at least, cooling radiation alone is sufficient to explain the extended Ly$\\alpha$ emission. However, often the constraints on the level of AGN or starburst activity are too poor to rule out the involvement of feedback energy completely.  In practice, it seems likely that a mixture of gravitational cooling and feedback processes are at play in LABs, each imparting an ambiguous observational signature in the Ly$\\alpha$ emission that is further muddied by observational limitations (namely angular and spectral resolution, and poor sensitivity to low surface-brightness features). In any case, it is important to note that {\\it both} the cooling and `heating' models require a reasonable reservoir of gas in the circumgalactic medium, so whatever the process dominating the Ly$\\alpha$ emission, LABs provide a rare opportunity to investigate the astrophysics at the galaxy/intergalactic medium interface close to the formation epoch of massive galaxies, $z\\approx2$--$3$. ",
        "conclusions": "\\subsection{A model of LAB formation: a combination of `velvet rope' and `bouncer' feedback?}  The observations imply that one of the main sources of Ly$\\alpha$ photons emitted by SSA22-LAB1 is either star formation or an AGN within the embedded ULIRG. These photons are escaping the galaxy and scattering in cold gas distributed in the CGM. Whatever the exact mechanism for Ly$\\alpha$ emission, there must be a large quantity of cold, neutral gas in the CGM. An open question therefore is the origin of the circumgalactic gas responsible for scattering/emitting the Ly$\\alpha$ photons and giving rise to the observed extended emission-line nebula. There are two options:  \\begin{enumerate} \\item The cold gas is comprised of pristine material containing dense clumps of H{\\sc i} being accreted onto the dark matter halo (Zheng et al.\\ 2011), potentially delivered by narrow streams within a hotter medium. Ly$\\alpha$ photons escaping the central galaxy illuminate this gas as they scatter in the clumpy, neutral medium. The morphology of the resulting LAB is therefore somewhat dependent on the geometry of the escaping photons and the covering factor of the incoming clumps, which could be modified as the inflow encounters radiation/outflows from the galaxy.  \\item Neutral gas is being ejected into the CGM via a galactic outflow, having already been accreted onto the galaxy. The outflow forms when shells of cold gas are driven outwards by pressurized bubbles inflated within the ISM via energy/momentum injection from stellar winds, supernovae and AGN. The outflow strongly resembles a rarefaction wave escaping the galaxy along paths of least resistance (e.g.,\\ perpendicular to the disc), accelerating as they emerge into the low-pressure CGM before fragmenting into cold clumps due to hydrodynamic instabilities and eventually decelerating in the gravitational potential of the halo (Dijkstra \\& Kramer 2012, Creasey et al.\\ 2013).\\end{enumerate}  \\noindent The key difference between (1) and (2) is the provenance and history of the cold gas dominanting the CGM: pristine still-cooling or previously cycled through a galaxy? The two scenarios describing the history of the CGM can be related to the concept of `velvet rope' and `bouncer' feedback\\footnote{terms attributed to Neal Katz}. In velvet rope feedback, gas never gets a chance to enter a galaxy where it can be assimilated into the ISM. In contrast, bouncer mode describes feedback that ejects gas from the galaxy back into the CGM.  Of course, these scenarios are not mutually exclusive. Gas must be delivered to growing galaxies at some point, and this is likely to be close to, or simultaneous with, the time the galaxies undergo their most rapid growth. Therefore it is possible that emergent shells comprised of enriched gas could also sweep-up cold, metal-free inflowing material that has not yet reached the disc (thus, potentially preventing it from being accreted: velvet rope feedback). In scenario 1, some of the inflowing gas might be blown back by an emergent wind, and the observations suggest that the bulk of the line-emitting gas is outflowing with respect to the central ULIRG (Fig.\\ 2). One problem with this scenario is that the covering factors of the streams are expected to be very small ($f\\sim1$--$3$ per cent; Faucher-Gigu\\`ere \\& Kere\\v{s}\\ 2011). Outflows are likely to emerge preferentially between these streams, allowing the cold inflowing material to continue to fuel the disc unarrested, while cold interstellar material is simultaneously ejected (bouncer feedback). However, the interaction of multi-phase galactic outflows and the narrow cooling streams predicted by models is poorly understood, the hydrodynamic interaction between the two is likely to be complicated, dependent on halo mass, geometry and radius. Moreover, the halo will still have a diffuse atmosphere, even if most of the mass is accreted along filamentary structure (e.g.,\\ Crain et al.\\ 2013).  While the cold flows model predicts that the inflowing cold gas should also be a source of Ly$\\alpha$ radiation from collisional excitation, the polarization observations show that a significant fraction of the Ly$\\alpha$ emission is not produced {\\it in situ}. However, in a scenario where Ly$\\alpha$ photons from the central source are scattering in H{\\sc i} clumps in these streams, one must consider the low covering factors predicted for the inflowing cold filaments. If the escaping Ly$\\alpha$ photons are uncorrelated with the geometry of the filaments (which could be a naive assumption if one considers that the escape must preferentially be in the directions of lowest density, i.e.,\\ potentially {\\it anti}-correlated with the filaments), then one would expect a fraction $(1-f)$ of the Ly$\\alpha$ photons {\\it not} to scatter. In effect, this drastically increases the required Ly$\\alpha$ luminosity of the central source. In contrast, gas entrained in a galaxy-scale wind can have covering factors approaching unity, which arguably would result in an observed Ly$\\alpha$ morphology similar to R2 and R3.  Francis et al.\\ (2013) argue that the CGM of Ly$\\alpha$ Blob B1 ($z=2.38$) could also be illuminated by scattered Ly$\\alpha$ radiation, either originating from a central source, or via fast shocks in the CGM. In the superwind scenario, the fast shock is formed when hot gas driven by the emergent wind from the central source collides with cold gas in the CGM. The detection of C~{\\sc iv} in the CGM of B1 is taken to be evidence for gas heating via fast shocks (Francis et al.\\ 2001), and Francis et al.\\ (2013) argue that the presence of carbon rules out a primordial origin for this material. Enriched cold gas in the vicinity of low-mass satellites within the halo could be the origin (these could also contribute Ly$\\alpha$ luminosity via ultraviolet leakage), as could metal-rich gas previously ejected from the central source. Still, a primordial origin for at least a fraction of the cold gas cannot be ruled out, and indeed should be expected.  \\subsection{Final comments}  It seems that a coherent picture for the SSA22-LAB1 system is at last emerging, although puzzles still remain. The SSA22-LAB1 system is clearly a group in the early stage of formation, and while the Ly$\\alpha$ `bridge' linking the northern LBG C11 with the main emission structure may well be a remnant cooling flow, the evidence suggests that a significant fraction of the Ly$\\alpha$ emission originates in the central ULIRG and scatters in H{\\sc i} clouds distributed in the CGM. The twist is that the central source does not emit Ly$\\alpha$ photons isotropically, indeed, we do not see ionizing photons from the source, let alone the escaping Ly$\\alpha$ radiation. Nevertheless, the new SCUBA-2 observations presented here allow us to observe the ultraviolet photons that are reprocessed by internal dust along our line-of-sight. Although we do not see the ultraviolet continuum source directly, abundant Ly$\\alpha$ photons (and presumably a strong flux of ionizing photons) must escape along an axis that is almost perpendicular to the line of sight. Viewed along the emergent axis from another part of the Universe we would almost certainly identify the source as a Ly$\\alpha$ halo containing an ultraviolet continuum source (as other LABs often do, e.g.,\\ Matsuda et al.\\ 2004).  SSA22-LAB1 offers a fascinating insight into the physics of massive galaxy formation: many of the ingredients of our latest models appear to be in action in this one system, and it clearly provides a unique opportunity to study the astrophysics of the galaxy/CGM interface. There are three key observations that would provide more important clues:  \\begin{enumerate}  \\item{Sensitive, high-resolution Ly$\\alpha$ imaging. This would more clearly reveal the morphology of the cold circumgalactic gas as traced by Ly$\\alpha$, providing the finer detail that might help distinguish an outflow with large covering factor from narrow filamentary streams. SSA22-LAB1 is an ideal target for the Multi Unit Spectroscopic Explorer (MUSE) IFU being commissioned on the Very Large Telescope.}  \\item{Morphology of the dust continuum. High resolution sub-mm imaging would more accurately pin-point the ULIRG within the Ly$\\alpha$ halo and potentially provide a clearer picture of the nature of the central luminous source, including the geometry of the obscuring material. The Atacama Large Millimetre Array (ALMA) is an obvious facility to perform this observation.}  \\item{Systemic redshift of the ULIRG. If the galaxy is responsible for a biconical outflow in the plane of the sky (as suggested by the relative velocities of the regions R2 and R3), then it should be very close in redshift space to the Ly$\\alpha$ peaks of R2 and R3. The ULIRG systemic redshift could be measured with millimetre observations of high-{\\it J} CO lines, which should trace the densest molecular gas at the sites of star formation within the ULIRG, and thus provide an accurate measurement of the systemic redshift relative to the extended Ly$\\alpha$. Again, this question could be addressed by ALMA.} \\end{enumerate}"
    },
    "1402/1402.4330_arXiv.txt": {
        "abstract": "{The problem of efficiently computing the orbital elements of a visual binary while still deriving confidence intervals with frequentist properties is treated. When formulated in terms of the Thiele-Innes elements, the known distribution of probability in Thiele-Innes space allows efficient grid-search plus Monte-Carlo-sampling schemes to be constructed for both the minimum-$\\!\\chi^{2}$ and the Bayesian approaches to parameter estimation. Numerical experiments with $10^{4}$ independent realizations of an observed orbit confirm that the $1-$ and $2\\sigma$ confidence and credibility intervals have coverage fractions close to their frequentist values. ",
        "introduction": "When error bars or confidence intervals are reported, the reader expects them to have their frequentist meaning. Thus, a 95\\% confidence interval is interpreted as implying a probability of 0.95 that the true result is enclosed by that interval. Similarly, the interval defined by $\\pm 1\\sigma$ error bars is expected to include the true answer with probability 0.683. However, this frequentist ideal is often not realized. This may be the result of observers misjudging the precision of their measurements or of large measurment errors occurring more frequently than expected for a normal distribution.  Such practical issues are absent when data analysis techniques are investigated with simulations, since precision can be exactly specified and measurement errors can be assigned with random gaussian variates, so that one might then expect a rigorous recovery of the frequentist ideal. But approximations can still compromise statistical rigour. For example, if a grid is required, confidence intervals might be affected if the grid is too coarse. In such cases, with increased computational resources, the limit as the grid steps $\\rightarrow 0$ can be closely approached and accurate results obtained.  Of more concern are approximations that compromise confidence intervals independently of any such limit. Two examples in the recent literature occur in {\\em hybrid} problems - i.e., non-linear problems with a subset of linear parameters. The first example is the code EXOFAST for analysing transit and radial velocity data for stars with orbiting planets (Eastman et al. 2013). These authors note that the convergence of their Markov Chain Monte Carlo (MCMC) parameter search is much faster if the exact solution for the linear parameters is introduced. However, the resulting uncertainties in the linear parameters are as much as 10 times smaller than when fitted non-linearly. Pending further research, these authors sensibly choose the inefficient option of treating all parameters as non-linear.  A similar but less extreme example arises when Bayesian estimation is applied to visual binaries (Lucy 2014; L14). When formulated in terms of the Thiele-Innes elements, the problem becomes linear in four of the seven elements. But when this linearity is exploited, coverage fractions (L14, Sect.5.5) indicate that the standard errors of the four linear elements are too small by factors of up to 2.1.  These examples pose a statistical challenge in the analysis of orbits: How can we benefit from partial linearity without losing the frequentist properties of confidence intervals? In this paper, this challenge is addressed in its visual binary context and for both frequentist and Bayesian procedures. ",
        "conclusions": "This paper has addressed a technical issue in the statistical analysis of orbits: how to achieve statistical rigour while taking advantage of the linearity of a subset of the orbital elements. The coverage experiments reported in Sects.3.5 and 4.5 demonstrate that statistical rigour is achieved for both minimum-$\\! \\chi^{2}$ and Bayesian estimation. Moreover, the reported timings show an inconsequential increase in the computational burden.  The key to this success is that at each grid point the distribution of probability in Thiele-Innes space is known. For minimum-$\\! \\chi^{2}$ estimation, this allows Monte Carlo sampling to be targeted (Sect.A.5) precisely on the constant $\\chi^{2}$ surface defining the desired confidence level. For Bayesian estimation, the known pdf allows Monte Carlo sampling in Thiele-Innes space to be concentrated (Sect.A.4) on the high probability domain enclosing the least-squares points $(\\hat{A},\\hat{B},\\hat{F},\\hat{G})$. Moreover, in each case, the random sampling does not require additional orbits to be computed.  The approach developed here is not specific to visual binaries or to the Thiele-Innes elements. In principle, an analogous procedure can be constructed for any partially linear estimation problem.   \\appendix"
    },
    "1402/1402.2336_arXiv.txt": {
        "abstract": "{We perform cosmological N-body simulations with non-Gaussian initial conditions generated from two independent fields. The dominant contribution to the perturbations comes from a purely Gaussian field, but we allow the second field to have local non-Gaussianity that need not be weak. This scenario allows us to adjust the relative importance of non-Gaussian contributions beyond the skewness, producing a scaling of the higher moments different from (and stronger than) the scaling in the usual single field local ansatz. We compare semi-analytic prescriptions for the non-Gaussian mass function, large scale halo bias, and stochastic bias against the simulation results. We discuss applications of this work to large scale structure measurements that can test a wider range of models for the primordial fluctuations than is usually explored. }  \\preprint{IGC-14/2-1}  \\begin{document} ",
        "introduction": "Recently, the {\\it Planck} satellite mission reported tight constraints on primordial non-Gaussianity from measurements of the Cosmic Microwave Background (CMB) \\cite{PlanckCollaboration2013}. They find the amplitude of the local, equilateral and orthogonal type bispectrum to be $\\fNL^{\\rm local}=2.7\\pm5.8$, $\\fNL^{\\rm equil}=-42\\pm75$, and $\\fNL^{\\rm ortho}=-25\\pm39$ respectively at $1\\sigma$ level. While these results show that the primordial fluctuations were remarkably Gaussian, they still leave room for interesting signatures of primordial physics to be found in statistics beyond the power spectrum: minimal, single field models predict non-Gaussianity that is about two orders of magnitude below these constraints. Non-Gaussianity is such an informative tool for inflationary physics that it is crucial to push observational bounds as far as possible. A number of forecasts have shown that the use of clustering data from future large scale structure surveys can constrain $\\fNL^{\\rm local}$ with $\\sigma(\\fNL)\\approx (1-10)$ \\cite{Pillepich2012, Cunha2010, Oguri2009, Giannantonio2012,Merloni:2012uf, Amendola:2012ys}. In addition, the galaxy bispectrum is a promising way to probe many shapes to the $\\fNL\\sim\\mathcal{O}(1-10)$ level \\cite{Jeong:2009vd,Baldauf:2010vn,Tasinato:2013vna}. The constraints from large scale structure (LSS) ---whether consistent with the CMB measurements or in tension --- will provide interesting and useful complementary results.  In the inflationary scenario, non-Gaussian signatures in the primordial density fluctuations depend on the details of interactions of the inflaton or other fields relevant for generating the perturbations; therefore, any detection (or the lack) of primordial non-Gaussianity tells us about the dynamics of those fields. Furthermore, if non-Gaussianity is detected, we expect to see patterns in the correlations that are consistent with perturbation theory (in contrast to apparently independent statistics at each order, for example). The simplest such pattern is in the relative amplitudes of higher order statistics, or, how the amplitudes of the correlation functions of the gravitational potential, $\\langle\\Phi^n(\\vec{x})\\rangle$, scale with order $n$. The relative scaling of higher moments falls into a fairly narrow range of behaviors in inflationary models \\cite{Barnaby2012}.  Thanks to the non-linearity of structure formation, the statistics of objects in the late universe contain information about the entire series of higher order moments of the initial density fluctuations. This is especially straightforward to see in number counts of galaxy clusters \\cite{Loverde2007}. While information about higher moments is clearly non-trivial to extract when non-Gaussianity is weak, some constraints with current data are already affected by assumptions about higher order moments. For example, \\cite{Shandera2013} reports constraints on primordial non-Gaussianity using a sample of 237 X-ray clusters from the ROSAT All-Sky Survey \\cite{Trumper1993}. Their analysis clearly suggests that it is important to take the scaling of moments into account when deriving constraints on $\\fNL$ from cluster number counts. The results presented in this paper will be useful for further studies in the same direction. Tighter constraints could likely be achieved with current data by combining the X-ray clusters with SZ detected clusters that have been separately used to constrain non-Gaussianity \\cite{Benson2013, Williamson2011, Mana:2013qba}. In addition, the eROSITA survey \\cite{Merloni:2012uf} and third generation surveys detecting clusters via the Sunyaev-Zel'dovich effect are expected to provide much larger samples of clusters in the next few years. It will be interesting to revisit the non-Gaussian analysis when that data becomes available. Other large scale structure statistics, including the power spectrum and bispectrum of galaxies, also contain contributions from higher order primordial statistics. Constraints on non-Gaussianity from those observables are still being developed, and will ultimately be very powerful (constraints obtained to date include \\cite{Afshordi:2008ru, Slosar2010, Xia:2010pe, Xia:2011hj, Ross:2012sx, Karagiannis:2013xea, Giannantonio:2013uqa, Ho:2013lda, Agarwal:2013qta}). To make full use of them it is important to understand the signatures of the full spectrum of non-Gaussian models predicted from inflation.  In this paper, we will use a slight variant of the popular \\textit{local} ansatz of primordial non-Gaussianity in order to generate two different scalings of the $n$-point functions and study how the scalings can be distinguished. For the usual local ansatz with \\textit{hierarchical} scaling, many groups have already performed simulations \\cite{Pillepich:2008ka, Grossi:2009an, Giannantonio:2009ak, Loverde2011, D'Amico:2010ta, Wagner2010} and found that the semi-analytic prescription of \\cite{Loverde2007} for the non-Gaussian mass function works well. Here we will also test the validity of the mass function proposed in \\cite{Barnaby2012} and used in \\cite{Shandera2013} for another scaling that is well motivated by inflationary models. In addition to looking at the mass function, we will also look at a signature of the shape (momentum dependence) of local type non-Gaussianity using the scale dependent halo bias \\cite{Dalal2008}, and at the stochastic bias on large scales \\cite{Tseliakhovich2010}. Subleading contributions to the bias and leading contributions to the stochastic bias are sensitive to non-Gaussianity beyond the skewness.  The rest of the paper is structured as follows: in Section \\ref{sec:Model} we introduce our two field model and the scaling of the higher moments. In Section \\ref{sec:Theory} we present the analytical calculations for the mass function of halos, bias, and stochasticity (with details relegated to the appendices). In Section \\ref{sec:sim} we describe our simulations, and then compare the output from simulations to the analytical calculations in Section \\ref{resultssection}. Finally, we conclude in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper, we have reported results for the mass function of massive halos and scale dependent bias from N-body simulations using a generalized two field model of primordial non-Gaussianity. The two field ansatz is used to generate initial conditions where the higher moments $\\mathcal{M}_{n}$, $n>3$ are more important than in the standard local ansatz, given the same value of $\\mathcal{M}_3$ as the single field local ansatz case. That this can be done is not mathematically surprising, but we have considered a range of scalings that are very natural from particle physics models of inflation. We have shown that in using large scale structure data to constrain the primordial fluctuations, assumptions about the scaling of higher moments should be explicitly stated. In addition, we have shown that in scenarios with more than one source for the density fluctuations it is important to allow independent bias coefficients {\\it for each source}.  Our simulations show that the Petrov expansion gives a good approximation to the non-Gaussian mass function when the amount of non-Gaussianity is small enough. The criteria for ``small enough\" depends on the dimensionless skewness $\\mathcal{M}_3$ {\\it and} the scaling of moments. We have verified that for the same level of skewness, the feeder scaling is more non-Gaussian than the hierarchical scaling by a straightforward comparison of simulation outputs of the two cases. For the parameter space we probed, we were able to show that the truncation error evaluation of the Petrov PDF, Eq.(\\ref{eq:pdf}), correlated well with the degree to which the truncated non-Gaussian mass function fitted the simulation results. This gives us further confidence in the use of the truncated Petrov (Edgeworth) mass function for various analyses such as that of \\cite{Shandera2012, Shandera2013} to put constraints on primordial non-Gaussianity using number counts of objects. We make progress towards calibrating the mass function formulae Eq.(\\ref{eq:hiermassfn}) and Eq.(\\ref{eq:feedmassfn}) for both scalings by using an extra parameter $f_2$, in addition to verifying that a reduced $\\delta_c$ is preferred, which is consistent with previous simulation studies of non-Gaussian mass functions with non-Gaussianity of local type and hierarchical scaling of higher moments. The effect this calibration might have on previous analysis of cluster constraints is illustrated in the bottom panels of Figure \\ref{mfnresults}.  Similarly, the peak-background-split calculations are good fits to the non-Gaussian bias results ($P_{hm}(k)/P_{mm}(k)$) from simulations---both for the hierarchical models (previously done a number of times) and for the feeder models. Moreover, we also presented results for large scale stochastic bias from our simulations which, in addition to the halo-matter bias, also depends on $P_{hh}(k)/P_{mm}(k)$. We found that, for two source scenarios, both the bias and stochastic bias calculations work well once we allow for different bias coefficients for the two independent fields. These bias coefficients, namely $b_\\phi$ and $b_\\psi$ in our notation, were measured by cross-correlating the halo density field with the linear density field contributions from $\\phi$ and $\\psi$ fields separately. The analytical calculations deviated from the simulation results for the stochastic bias only for halo samples whose values of stochasticity were inconsistent with unity in the Gaussian simulations. Therefore, for these halo samples, we expect the need to account for other contributions to the stochasticity that have to be taken into account even in the Gaussian case. While we have not investigated this issue in detail, additional studies would be worthwhile since a better understanding of stochasticity is useful for cosmological applications of galaxy surveys.   \\paragraph{Acknowledgment:} This work is supported by the National Aeronautics and Space Administration under Grant No. NNX12AC99G issued through the Astrophysics Theory Program. This work used the Extreme Science and Engineering Discovery Environment (XSEDE), which is supported by National Science Foundation grant number OCI-1053575. The simulations for the project were performed using computing resources at the Texas Advanced Computing Center (TACC), The University of Texas at Austin through XSEDE grant AST130053 and at the Penn State Research Computing and Cyberinfrastructure (RCC). S. A. would like to thank Xinghai Zhao for suggestions on setting up Gadget simulations.  \\appendix"
    },
    "1402/1402.0043_arXiv.txt": {
        "abstract": "Direction-sensitive WIMP dark matter searches may help overcome the challenges faced by direct dark matter detection experiments. In particular, directional detectors should be able to clearly differentiate a dark matter signal from background sources. We are developing a Directional Dark Matter Detector (D$^3$) based on a gas Time Projection Chamber (TPC) using Gas Electron Multipliers (GEMs) for charge amplification and pixel electronics for readout. This approach allows the three-dimensional reconstruction of nuclear recoils in a room-temperature detector with low energy threshold and low noise. We present an overview of our past and ongoing work developing this technology, including the performance measurements of small prototypes, as well as our planned future work constructing a m$^3$-scale detector to clearly determine whether the signals seen by DAMA, CoGeNT, and CRESST-II are due to low-mass WIMPs or background. ",
        "introduction": "Our group is investigating the feasibility of micro-pattern gas detectors for use in direction-sensitive searches for weakly intreating massive particles (WIMPs). We are specifically evaluating the use of Gas Electron Multipliers (GEMs) and pixel electronics as readout technology for a gas Time Projection Chamber (TPC). The \\dcube readout system is one of four candidate technologies for DRIFT-HD, a planned large-scale directional detector to be built as part of the DRIFT (Directional Recoil Identification From Tracks) experiment.  \\begin{figure} \\centering \\includegraphics[width=0.7\\linewidth]{dailyosc2.pdf} \\caption{The rotation of Earth as it moves through the WIMP wind causes a daily shift in the incoming particle direction as measured in the lab frame \\cite{cygnus}.} \\label{dailyosc} \\end{figure}  Directionally-sensitive dark matter detectors are of interest due to their potential to provide unambiguous evidence for the existence of WIMPs. The earth's motion with respect to the galactic rest frame should produce a ``WIMP wind\" through the planet originating from the direction of the constellation Cygnus. The rotation of the earth as it moves through this wind would cause a daily 90$^{\\circ}$ shift in the incoming particle direction as seen by a detector on the surface. This principle is illustrated in Figure \\ref{dailyosc}. There is no known background with a similar signature, making this a very promising detection method. It is also a very large effect, with some models predicting that the signal could be distinguished from background with only $\\mathcal{O}$(10) events \\cite{copi}. ",
        "conclusions": "The results of the \\dcubem prototype show that a TPC using GEMs and pixel readout can achieve precise, 3D directional detection of nuclear recoils. The detector can be run stably for extended periods, and the GEMs allow for high gain with good resolution, which should enable accurate head-tail identification for recoils. Overall, the results are promising for the application of this detector technology to the search for WIMPs.  \\begin{figure} \\centering \\includegraphics[width=0.7\\linewidth]{dcubemilli.png} \\caption{Top-down view of \\dcubemilli design, which implements four unit cells inside a common field cage.} \\label{dcubemilli} \\end{figure}  In the immediate future, we will begin work on a larger prototype detector, \\dcubemilli (depicted in Figure \\ref{dcubemilli}). This detector will be much larger than the current \\dcubem, with approximately 10 liters of active volume. It will use the new FE-I4 pixel chip along with 10 $\\times$ 10 cm GEMs from CERN, as well as an upgraded DAQ system.  We have also recently joined the DRIFT Experiment, the world's leading directional dark matter search effort. DRIFT is the only experiment running a directionally-sensitive detector on the m$^3$ scale, and they are currently working on plans for DRIFT-III, a 24 m$^3$ detector with 4 kg of target mass. We believe that combining our readout technology with DRIFT's experience and expertise will give us a great advantage in building a competitive directionally-sensitive WIMP detector."
    },
    "1402/1402.6703.txt": {
        "abstract": "Past studies have identified a spatially extended excess of $\\sim$1-3 GeV gamma rays from the region surrounding the Galactic Center, consistent with the emission expected from annihilating dark matter. We revisit and scrutinize this signal with the intention of further constraining its characteristics and origin. By applying cuts to the \\textit{Fermi} event parameter CTBCORE, we suppress the tails of the point spread function and generate high resolution gamma-ray maps, enabling us to more easily separate the various gamma-ray components. Within these maps, we find the GeV excess to be robust and highly statistically significant, with a spectrum, angular distribution, and overall normalization that is in good agreement with that predicted by simple annihilating dark matter models. For example, the signal is very well fit by a 36-51 GeV dark matter particle annihilating to $b\\bar{b}$ with an annihilation cross section of $\\sigma v = (1-3)\\times 10^{-26}$ cm$^3$/s (normalized to a local dark matter density of 0.4 GeV/cm$^3$). Furthermore, we confirm that the angular distribution of the excess is approximately spherically symmetric and centered around the dynamical center of the Milky Way (within $\\sim$$0.05^{\\circ}$ of Sgr A$^*$), showing no sign of elongation along the Galactic Plane. The signal is observed to extend to at least $\\simeq10^{\\circ}$ from the Galactic Center, disfavoring the possibility that this emission originates from millisecond pulsars. ",
        "introduction": "\\label{intro}  Weakly interacting massive particles (WIMPs) are a leading class of candidates for the dark matter of our universe. If the dark matter consists of such particles, then their annihilations are predicted to produce potentially observable fluxes of energetic particles, including gamma rays, cosmic rays, and neutrinos. Of particular interest are gamma rays from the region of the Galactic Center which, due to its proximity and high dark matter density, is expected to be the brightest source of dark matter annihilation products on the sky, hundreds of times brighter than the most promising dwarf spheroidal galaxies.  Over the past few years, several groups analyzing data from the \\textit{Fermi} Gamma-Ray Space Telescope have reported the detection of a gamma-ray signal from the inner few degrees around the Galactic Center (corresponding to a region several hundred parsecs in radius), with a spectrum and angular distribution compatible with that anticipated from annihilating dark matter particles~\\cite{Goodenough:2009gk,Hooper:2010mq,Boyarsky:2010dr,Hooper:2011ti,Abazajian:2012pn,Gordon:2013vta,Abazajian:2014fta}. More recently, this signal was shown to also be present throughout the larger Inner Galaxy region, extending kiloparsecs from the center of the Milky Way~\\cite{Hooper:2013rwa,Huang:2013pda}. While the spectrum and morphology of the Galactic Center and Inner Galaxy signals have been shown to be compatible with that predicted from the annihilations of an approximately 30-40 GeV WIMP annihilating to quarks (or a $\\sim$7-10 GeV WIMP annihilating significantly to tau leptons), other explanations have also been proposed. In particular, it has been argued that if our galaxy's central stellar cluster contains several thousand unresolved millisecond pulsars, they might be able to account for the emission observed from the Galactic Center~\\cite{Hooper:2010mq,Abazajian:2010zy,Hooper:2011ti,Abazajian:2012pn,Gordon:2013vta,Abazajian:2014fta}. The realization that this signal extends well beyond the boundaries of the central stellar cluster~\\cite{Hooper:2013rwa,Huang:2013pda} disfavors such interpretations, however.  In particular, pulsar population models capable of producing the observed emission from the Inner Galaxy invariably predict that \\textit{Fermi} should have resolved a much greater number of such objects. Accounting for this constraint, Ref.~\\cite{Hooper:2013nhl} concluded that no more than $\\sim$5-10\\% of the anomalous gamma-ray emission from the Inner Galaxy can originate from pulsars. Furthermore, while it has been suggested that the Galactic Center signal might result from cosmic-ray interactions with gas~\\cite{Hooper:2010mq,Hooper:2011ti,Abazajian:2012pn,Gordon:2013vta}, the analyses of Refs.~\\cite{Linden:2012iv} and~\\cite{Macias:2013vya} find that measured distributions of gas provide a poor fit to the morphology of the observed signal.  It also appears implausible that such processes could account for the more spatially extended emission observed from throughout the Inner Galaxy.     In this study, we revisit the anomalous gamma-ray emission from the Galactic Center and the Inner Galaxy regions and scrutinize the \\textit{Fermi} data in an effort to constrain and characterize this signal more definitively, with the ultimate goal being to confidently determine its origin. One way in which we expand upon previous work is by selecting photons based on the value of the \\textit{Fermi} event parameter CTBCORE. Through the application of this cut, we select only those events with more reliable directional reconstruction, allowing us to better separate the various gamma-ray components, and to better limit the degree to which emission from the Galactic Disk leaks into the regions studied in our Inner Galaxy analysis.  We produce a new and robust determination of the spectrum and morphology of the Inner Galaxy and the Galactic Center signals. We go on to apply a number of tests to this data, and determine that the anomalous emission in question agrees well with that predicted from the annihilations of a 36-51 GeV WIMP annihilating mostly to $b$ quarks (or a somewhat lower mass WIMP if its annihilations proceed to first or second generation quarks). Our results now appear to disfavor the previously considered 7-10 GeV mass window in which the dark matter annihilates significantly to tau leptons~\\cite{Hooper:2010mq,Hooper:2011ti,Gordon:2013vta,Hooper:2013rwa,Abazajian:2014fta} (the analysis of Ref.~\\cite{Gordon:2013vta} also disfavored this scenario). The morphology of the signal is consistent with spherical symmetry, and strongly disfavors any significant elongation along the Galactic Plane. The emission decreases with the distance to the Galactic Center at a rate consistent with a dark matter halo profile which scales as $\\rho \\propto r^{-\\gamma}$, with $\\gamma \\approx 1.1-1.3$. The signal can be identified out to angles of $\\simeq 10^{\\circ}$ from the Galactic Center, beyond which systematic uncertainties related to the Galactic diffuse model become significant. The annihilation cross section required to normalize the observed signal is $\\sigma v \\sim 10^{-26}$ cm$^3$/s, in good agreement with that predicted for dark matter in the form of a simple thermal relic.   The remainder of this article is structured as follows. In the following section, we review the calculation of the spectrum and angular distribution of gamma rays predicted from annihilating dark matter.  In Sec.~\\ref{ctbcore}, we describe the event selection used in our analysis, including the application of cuts on the \\textit{Fermi} event parameter CTBCORE. In Secs.~\\ref{inner} and~\\ref{center}, we describe our analyses of the Inner Galaxy and Galactic Center regions, respectively. In each of these analyses, we observe a significant gamma-ray excess, with a spectrum and morphology in good agreement with that predicted from annihilating dark matter. We further investigate the angular distribution of this emission in Sec.~\\ref{morphology}, and discuss the dark matter interpretation of this signal in Sec.~\\ref{darkmatter}. In Sec.~\\ref{discussion} we discuss the implications of these observations, and offer predictions for other upcoming observations. Finally, we summarize our results and conclusions in Sec.~\\ref{summary}. In the paper's appendices, we include supplemental material intended for those interested in further details of our analysis.  \\bigskip     \\begin{figure*}[t!] \\includegraphics[width=6.9in]{profiles.pdf} \\caption{Left frame: The dark matter density as a function of the distance to the Galactic Center, for several halo profiles, each normalized such that $\\rho=0.4$ GeV/cm$^3$ at $r=8.5$ kpc. Right frame: The line-of-sight integral of the density squared, as defined in Eq.~\\ref{J}, for the same set of halo profiles, as a function of the angular distance from the Galactic Center, $\\psi$.} \\label{profiles} \\end{figure*} ",
        "conclusions": "\\label{summary}  In this study, we have revisited and scrutinized the gamma-ray emission from the central regions of the Milky Way, as measured by the \\textit{Fermi} Gamma-Ray Space Telescope. In doing so, we have confirmed a robust and highly statistically significant excess, with a spectrum and angular distribution that is in excellent agreement with that expected from annihilating dark matter. The signal is distributed with approximate spherical symmetry around the Galactic Center, with a flux that falls off as $F_{\\gamma} \\propto r^{-(2.2-2.6)}$, implying a dark matter distribution of $\\rho \\propto r^{-\\gamma}$, with $\\gamma \\simeq 1.1-1.3$. The spectrum of the excess peaks at $\\sim$1-3 GeV, and is well fit by 36-51 GeV dark matter particles annihilating to $b\\bar{b}$. The annihilation cross section required to normalize this signal is $\\sigma v = (1.9-2.8) \\times 10^{-26}$ cm$^3$/s (for a local dark matter density of 0.4 GeV/cm$^3$), in good agreement with the value predicted for a simple thermal relic. In particular, a dark matter particle with this cross section will freeze-out of thermal equilibrium in the early universe to yield an abundance approximately equal to the measured cosmological dark matter density (for the range of masses and cross sections favored for other annihilation channels, see Sec.~\\ref{darkmatter}).  In addition to carrying out two different analyses (as described in Secs.~\\ref{inner} and~\\ref{center}), subject to different systematic uncertainties, we have applied a number of tests to our results in order to more stringently determine whether the characteristics of the observed excess are in fact robust and consistent with the signal predicted from annihilating dark matter. These tests uniformly confirm that the signal is present throughout the Galactic Center and Inner Galaxy (extending out to angles of at least $10^{\\circ}$ from the Galactic Center), without discernible spectral variation or significant departures from spherical symmetry. No known, anticipated, or proposed astrophysical diffuse emission mechanisms can account for this excess. And while a population of several thousand millisecond pulsars could have plausibly been responsible for much of the anomalous emission observed from within the innermost $\\sim1^{\\circ}-2^{\\circ}$ around the Galactic Center, the extension of this signal into regions well beyond the confines of the central stellar cluster strongly disfavors such objects as the primary source of this signal. In light of these considerations, we consider annihilating dark matter particles to be the leading explanation for the origin of this signal, with potentially profound implications for cosmology and particle physics.   \\bigskip  \\bigskip    {\\it Acknowledgements}: We would like to thank Keith Bechtol, Eric Charles and Alex Drlica-Wagner for their help with the Fermi-LAT likelihood analysis, Oscar Macias-Ramirez and Farhad Yusef-Zadeh for providing the 20 cm templates, and Simona Murgia, Jesse Thaler and Neal Weiner for helpful discussions. We particularly thank Francesca Calore, Ilias Cholis and Christoph Weniger for providing an independent cross-check of some of our sphericity results. We acknowledge the University of Chicago Research Computing Center for providing support for this work. DPF is supported in part by the NASA Fermi Guest Investigator Program, DH is supported by the Department of Energy, and TL is supported by NASA through Einstein Postdoctoral Award Number PF3-140110. TRS is supported by the U.S. Department of Energy under cooperative research agreement Contract Number DE-FG02-05ER41360."
    },
    "1402/1402.6665_arXiv.txt": {
        "abstract": "We present an analysis of three archival Chandra observations of the black hole V4641~Sgr, performed during a decline into quiescence.  The last two observations in the sequence can be modeled with a simple power-law.  The first spectrum, however, is remarkably similar to spectra observed in Seyfert-2 active galactic nuclei, which arise through a combination of obscuration and reflection from distant material.  This spectrum of V4641~Sgr can be fit extremely well with a model including partial-covering absorption and distant reflection. This model recovers a $\\Gamma \\simeq 2.0$ power-law incident spectrum, typical of black holes at low Eddington fractions.  The implied geometry is plausible in a high-mass X-ray binary like V4641~Sgr, and may be as compelling as explanations invoking Doppler-split line pairs in a jet, and/or unusual Comptonization.  We discuss potential implications and means of testing these models. ",
        "introduction": "The nature of X-ray emission at low Eddington fractions remains uncertain.  It is particularly interesting to explore connections between accretion inflows and jet outflows in X-ray binaries at $L\\simeq 10^{-4}~{\\rm L}_{Edd}$, as this regime may provide insights accross the mass-scale into the much larger class of low-luminosity active galactic nuclei (LLAGN) which exist at comparably low Eddington fractions.  Recent work on V404 Cyg, for instance, has shown that jet production can occur over many orders of magnitude (see e.g. Gallo, Fender, \\& Pooley\\ 2003).  One particularly well-studied jet-producing X-ray binary is SS 433. SS 433 is unique in that its jets can be resolved in different wavelengths, removing ambiguities in how spectral signatures map onto different parts of the accretion flow.  In particular, strong X-ray emission lines from the extended jets strongly imply an outflow power that exceeds the radiative Eddington limit for reasonable compact object masses (see, e.g., Marshall et al.\\ 2013).  Of course, the mechanical luminosity is not governed by the radiative Eddington limit, but since the power in the jets is likely generated in part via accretion, it is possible that the mass accretion rate in SS 433 is super-Eddington.  It has recently been suggested that V4641 Sgr may be a precessing micro-blazar that launches extreme jets, based on an unusual {\\it Chandra} X-ray spectrum obtained at low luminosity (Gallo, Plotkin, \\& Jonker 2014).  The spectrum shows a strange hump or break, near to and above 4 keV.  It may require a combination of relativistically--shifted emission lines from jets, as in the case of SS 433, and/or unusual continuum emission also tied to the jets.  It is notable that the feature(s) detected in V4641 Sgr are far stronger than lines detected in SS 433.  However, in any high-mass X-ray binary (HMXB), the massive companion wind may serve to complicate inferences regarding the accretion flow. Such winds can be clumpy, and can give rise to aperiodic obscuration that may fully or partially cover the central engine (the well-known dips in Cygnus X-1 are one prominent example of clumps; see e.g. Hanke et al.\\ 2009 ).  Such obscuration has been observed on multiple occasions in the spectrum of V4641 Sgr in both optical and X-ray wavelengths (Maitra \\& Bailyn\\ 2006; Charles et al.\\ 1999; Revnivtsev et al.\\ 2002b). Even in the absence of clumps, shocks and ionization fronts can lead to an inhomogeneous medium (see, e.g., Watanabe et al.\\ 2006).  When the direct continuum is wholly or partly blocked in a HMXB, the equivalent width of emission lines (arising via irradiation of the wind) can become very large (even 1--2 keV or more; see White, Nagase, \\& Parmar 1995).  The full complexity of HMXB winds has become clear in the era of advanced CCDs and dispersive X-ray spectrometers (see, e.g., Torrejon et al.\\ 2010).  In Vela X-1, which harbors a B-type star (albeit one with a stronger wind than V4641~Sgr), excellent X-ray spectra and complex wind models have been compared; a wind with multiple zones, shocks, and columns that can approach $N_{H} \\geq 10^{23}~{\\rm cm}^{-2}$ is clearly indicated (Watanabe et al. 2006).  V4641 Sgr is highly unusual even among HMXBs, and unique in other respects.  The system harbors a black hole (${\\rm M}_{BH} = 9.6^{+2.1}_{-0.9}~{\\rm M}_{\\odot}$) accreting from a high-mass B9 III-type (${\\rm M}_{C} = 6.5^{+1.6}_{-1.0}~{\\rm M}_{\\odot}$) companion star (Orosz et al.\\ 2001, although MacDonald et al. 2014 find slightly lower component masses: ${\\rm M}_{BH}=6.4\\pm0.6~{\\rm M}_{\\odot}$ and ${\\rm M}_{C}=2.9\\pm0.4~{\\rm M}_{\\odot}$).  However, the inclination angle of the superluminal jet detected in 1999 by Hjellming et al. (2000) is likely $i_{jet} \\leq 10^{\\circ}$ , whereas the inclination of the binary system is likely much higher, $60^{\\circ} < i_{binary} < 70^{\\circ}$ (Orosz et al.\\ 2001).  It is not clear how this geometry might affect observed spectra.   In this work, we examine {\\it Chandra} observations of V4641~Sgr at low Eddington fractions, originally analyzed by Gallo, Plotkin, \\& Jonker (2014).  We suggest an alternative explanation for the SS~433-like features that may be present in the spectrum.  The observations and data reduction are described in Secton 2.  In Section 3, we report the results of our analysis.  Finally we discuss our results in section 4. ",
        "conclusions": "We have reanalyzed three archival {\\it Chandra} spectra of V4641~Sgr, obtained as the source declined into a quiescent state.  The first spectrum in the sequence (observation 4451; see Table 1 and Figures 1 and 2) is of particular interest, as it shows strong deviations from a simple power-law.  This spectrum is very similar to those commonly obtained from Seyfert-2 AGN (e.g. Turner et al.\\ 1997; Matt et al.\\ 2000; Risaliti, Elvis, \\& Nicastro 2002; Bianchi et al.\\ 2009).  When fit with a model based on Seyfert-2 spectra -- one including partial covering absorption and distant reflection -- an excellent fit is obtained, and a typical incident spectrum is recovered.  In this section, we examine ways by which this spectrum might arise in V4641~Sgr, implications for accretion flows onto black holes, and means of testing this model.  As briefly mentioned above, V4641~Sgr is not a standard X-ray binary. It is the only transient high mass X-ray binary with a dynamically--confirmed black hole primary (Orosz et al.\\ 2001).  Our expectations of high mass X-ray binaries harboring black holes are built upon persistent (but variable) sources such as Cygnus X-1, LMC X-1, and LMC X-3.  The phenomena possible when black holes enter and return to quiescence in the presence of a high-mass companion are unexplored apart from V4641~Sgr.  Although the companion wind in V4641~Sgr is likely to be weaker than in a source like Vela X-1, accretion-induced changes in photoionization and shocks in a companion wind may cause inhomogeneities within it, and could lead to a complex absorption geometry (see, e.g., Watanabe et al.\\ 2006).  Of course, if the wind is clumpy, then a complex absorption geometry would also be likely.  Recently MacDonald et al. (2014) revisited the orbital parameters of V4641 Sgr by examining photometric data obtained over the last decade. They found that the companion star to V4641 Sgr exists in two states: a ``passive'' state in which the secondary exhibits very little intrinsic variability and the variations in the optical light curve are caused by the orbital modulation, and an ``active'' state, in which variability in the light curve is dominated by the intrinsic variability in the secondary.  Based on MacDonald et al.\\ (2014) and prior work by Orosz et al.\\ (2001), we find that observation 4451 occurs during the active state of the secondary, and observations 4452 and 4453 occur in the passive state.  That observations 4452 and 4453 are so different from observation 4451 suggests that the peculiar spectrum may be partially caused by the increased/changed wind geometry etc. from the secondary.  An active X-ray state combined with an active companion (wind) state could cause the shocks and fronts that would give rise to the observed spectrum, aided or unaided by clumps.  In a study of Seyfert galaxies, Akylas \\& Georgantopoulos (2009) found that the lower luminosity sources more often had unabsorbed nuclei, while the higher luminosity sources often had internal hydrogen column densities as high as $\\approx10^{24}~{\\rm cm}^{-2}$.  The unabsorbed sources had Eddington-scaled luminosities comparable to observations 4452 and 4453, while observation 4451 had an Eddington-scaled luminosity closer to that of the obscured sources.  That fact that the absorption behavior observed in V4641~Sgr seems to mirror Seyfert behavior may simply be coincidence.  However, it could represent a physical similarity.  One possibility may be a clumpy wind from the accretion flow operates at the Eddington fraction of observation 4451, but disappears at lower luminosity.  Elitzur \\& Shlosman (2006) developed such a model for AGN.  V4641~Sgr is also unique in the degree by which the inner accetion disk appears to be misaligned with the binary system ($i \\leq 10^{\\circ}$ versus $i = 60-70^{\\circ}$; Orosz et al.\\ 2001).  Depending on the vertical profile of the disk, the nature of the inner accretion flow, and how these geometries vary with the mass accretion rate, it is possible that the outer disk may act like the torus in AGN, and serve to partly or nearly entirely obscure the inner disk.  This is another means by which a Seyfert-2--like geometry and spectrum might arise in V4641~Sgr; it is independent of the companion wind, but of course the two could act in concert.   In order to model the spectrum of observation 4451, prior efforts developed some jet-focused models, including extremely strong Doppler-split lines from SS-433--like jets and/or peculiar Comptonization (Gallo, Plotkin, \\& Jonker 2014).  If the 4--8 keV spectrum is due to a set of lines, the lines are far stronger than those observed in SS~433 (e.g. Marshall et al.\\ 2013), even though a far higher Eddington fraction would be estimated from SS~433 if it were viewed down its axis.  Indeed, the jets in V4641~Sgr would have a bulk Lorentz factor of 15--45 (Gallo, Plotkin, \\& Jonker 2014).  While Lorentz factors this large are inferred in jets from supermassive black holes, they are considerably higher than the Lorentz factors typically observed in jets originating from other stellar mass black hole systems in the hard or quiescent states (see for example Fender, Belloni, \\& Gallo 2004).  Moreover, the continuum implied in such fits is unusually hard, and inconsistent with standard Comptonization ( $\\Gamma =$0.93--1.07), whereas our model recovers a more typical continuum.  It is possible that precessing jets have been observed in V4641~Sgr. However, in view of the qualitative and quantitative similarities between observation 4451 and well-known Seyfert-2 spectra, in view of the likely complexity of the environment in this transient, mis-aligned, high-mass X-ray binary, and in view of the fact that a canonical continuum can be recovered with an absorption plus reflection model, we suggest that the latter is also a compelling interpretation of the data.   One way to constrain the nature of the absorption is to examine the orbital geometry of the system at the times of our observations. In Cygnus X-1 at least, obscuration does appear to be phase--dependent. Any phase dependence may help to distinguish how the spectrum obtained in observation 4451 arose, so we calculated the phase ($\\phi$) of all three observations using orbital period and photometric $T_{0}$ calculated by Orosz et al. (2001; photometric $T_{0}$ is defined by Orosz et al. as the time of the superior conjunction of the secondary).  The phases are listed in Table~\\ref{tab:par} along with the uncertainties estimated from uncertainties in their phase zero-point and orbital period.  MacDonald et al.\\ (2014) estimated a more recent $T_{0}$ using the same orbital period from Orosz et al.\\ (2001).  Folding our observations on that $T_{0}$, the phase is comparable to that found by Orosz et al. Thus our choices of $T_{0}$ and $P_{\\rm orb}$ (and, in turn, the phases we compute using them) are consistent with more recently performed studies of the secondary.  If we then accept these phase determinations, we find that observation 4451 occurs near phase 0.0 (i.e. when the black hole is closest to us).  This nominally implies that the absorption may not be tied to an enhancement in the wind that is constant in binary phase, and suggests that the absorption is more closely tied to increased activity in the secondary and/or the inner accretion flow, and/or related effects on the outer accretion disk.  In summary, the unusual spectrum revealed in V4641~Sgr is similar to spectra obtained in Seyfert-2 AGN, and we suggest that a variety of plausible scenarios -- some tied to the HMXB nature of V4641~Sgr, others to its odd system parameters -- could produce the geometry consistent with a Seyfert-2--like spectral model.  Other origins for the spectrum -- including X-ray line emission from extreme jets and/or Comptonization -- cannot be ruled out, but we argue that partial-covering absorption and distant reflection provide a familiar explanation.  Future observations that sample a declining X-ray flux trend in V4641~Sgr, and/or the orbital phase, may be able to better reveal the nature of the spectrum.  The authors acknowledge helpful and supportive conversations with Elena Gallo, Richard Plotkin, and Peter Jonker.  We also thank the anonymous referee for a helpful review of this work."
    },
    "1402/1402.4456_arXiv.txt": {
        "abstract": "We use thermal radiometry and visible photometry to constrain the size, shape, and albedo of the large Kuiper belt object Haumea. The correlation between the visible and thermal photometry demonstrates that Haumea's high amplitude and quickly varying optical light curve is indeed due to Haumea's extreme shape, rather than large scale albedo variations.  However, the well-sampled high precision visible data we present does require longitudinal surface heterogeneity to account for the shape of lightcurve. The thermal emission from Haumea is consistent with the expected Jacobi ellipsoid shape of a rapidly rotating body in hydrostatic equilibrium. The best Jacobi ellipsoid fit to the visible photometry implies a triaxial ellipsoid with axes of length 1920 x 1540 x 990 km and density 2.6 g cm$^{-3}$, as found by \\citet{Lellouch10}. While the thermal and visible data cannot uniquely constrain the full non-spherical shape of Haumea, the match between the predicted and measured thermal flux for a dense Jacobi ellipsoid suggests that Haumea is indeed one of the densest objects in the Kuiper belt. ",
        "introduction": "\\label{sec:intro} Haumea, one of the largest bodies in the Kuiper Belt, is also one of the most intriguing objects in this distant population.  Its rapid rotation rate, multiple satellites, and dynamically-related family members all suggest an early giant impact \\citep{Brown07}.  Its surface spectrum reveals a nearly pure water ice surface \\citep{Barkume06,Trujillo07,Merlin07,Dumas11}, with constraints on other organic compounds with upper bounds $<$ 8\\% \\citep{Pinilla09}.  \\citet{Lacerda08} also find evidence for a dark spot on one side of the rotating body, make the surface albedo non-uniform.  Even more interesting is that shape modeling has suggested a density higher than nearly anything else known in the Kuiper belt and consistent with a body almost thoroughly dominated by rock \\citep{Rabinowitz06,Lacerda08,Lellouch10}. \\citet{Lacerda07} concluded a density of 2.551 g cm$^{-3}$ and follow up work found consistent values between 2.55 and 2.59 g cm$^{-3}$, depending on the model used.  Such a rocky body with an icy exterior could be a product of initial differentiation before giant impact and subsequent removal of a significant amount of the icy mantle.  \\citet{Leinhardt10} demonstrate that such an impact is possible in a graze and merge collision between equal size bodies, and are able to reproduce the properties of the Haumea family system.  Much of our attempt at understanding the history of Haumea relies on the estimate of the high density of the body.  Haumea's large-amplitude light curve and rapid rotation have been used to infer an elongated shape for the body. Assuming that the rotation axis lies in the same plane as the plane of the satellites, the amplitude of the light curve then gives a ratio of the surface areas along the major and minor axes of the bodies. If the further assumption is made that Haumea is large enough to be in hydrostatic equilibrium, the full shape can be uniquely inferred to be a Jacobi ellipsoid with fixed ratios of the three axes.  From this shape and from the known rotation velocity, the density is precisely determined.  Finally, with the mass of Haumea known from the dynamics of the two satellites \\citep{Ragozzine09}, the full size and shape of Haumea is known \\citep{Rabinowitz06,Lacerda08,Lellouch10}.  The major assumption in this chain of reasoning is that Haumea is a figure of equilibrium.  While it is true that a strengthless body can instantaneously have shapes very different from figures of equilibrium \\citep{Holsapple07}, long-term deformation at the pressures obtained in a body this size should lead to essentially fluid behavior, at least at depth. Indeed, any non-spinning body large enough to become round due to self-gravity has attained the appropriate figure of equilibrium. While the size at which this rounding occurs in the outer solar system is not well known, the asteroid Ceres, with a diameter of 900 km is essentially round \\citep{Millis87,Thomas05}, while among the icy satellites, everything the size of Mimas and larger ($\\sim$400 km) is essentially round. It thus seems reasonable to assume that a non-spinning Haumea, with a diameter of $\\sim$1240 km \\citep{Lellouch13,Fornasier13} would be round, thus a rapidly-spinning Haumea should likewise have a shape close to that of the Jacobi ellipsoid defined by Haumea's density and spin rate.  Nonetheless, given the importance of understanding the interior structure of Haumea and the unusually high density inferred from these assumptions, we find it important to attempt independent size and density measurements of this object. Here we use unpublished photometric data from the Hubble Space Telescope (HST) to determine a best-fit Jacobi ellipsoid model. We then compare the predicted thermal flux from this best-fit Jacobi ellipsoid to thermal flux measured from the Spitzer Space Telescope(first presented in \\citet{Stansberry08}) and consider these constraints on the size and shape of the body. ",
        "conclusions": "\\label{sec:disc} The thermal and photometric light curves of Haumea are consistent with the assumption that Haumea is a fluidly relaxed, rapidly rotating Jacobi ellipsoid with a density of 2.6 g cm$^{-3}$ and minor albedo variation across its surface. Although we allow the albedo to change longitudinally, we have demonstrated that the reason for the double-peaked lightcurve of Haumea is in fact a shape effect. This work presents new data and an informed Hapke model that agree with the findings of previous authors \\citep{Lellouch10,Lacerda08,Rabinowitz06}.  The lightcurve in several colors \\citep{Lacerda08,Lacerda09} indicates that the albedo variation is concentrated in a large spot on one side of the body, an argument supported by our allowed albedo variation.  The precise geometric albedo presented here is less than that of \\citet{Fornasier13}, but with a small variation in either the single-scattering albedo or the asymmetry parameter, we easily find agreement between the two values.  We do not focus on the absolute value of the albedo here, as it is closely tied to the unknown Hapke parameters for the surface.  The only affect this has on the thermal fit is to change the best fit beaming parameter, which does not differ by more than 1-$\\sigma$ from the reported value.  The important point here, however, is that our precise visible data set agrees with the dark spot proposed by \\citet{Lacerda08}.  There are still a number of questions remaining regarding this KBO.  The temperature of the dark spot is still uncertain and our data are not precise enough to constrain it, although the lower albedo we use to match the visible data agree well with a warmer region.  Another assumption used here that could be disputed is the rotation axis of the body as perpendicular to our line of sight.  However, assuming that the majority of the depth of the light curve is from the major axes of the body, the pole position will only affect the size of the third dimension of the body and the thermal beaming factor, which become somewhat degenerate when fitting the thermal data anyway.  If the body were not mostly edge-on, it would be difficult to explain the particular surface patterning needed to recreate the visible lightcurve.  While it is encouraging that a dense Jacobi ellipsoid fits both data sets, it is important to point out that we can only obtain a unique solution under the assumption that the object has a prescribed shape. If Haumea is modeled as an arbitrary triaxial ellipsoid rather than a Jacobi ellipsoid, large families of solutions are possible. To first order, the photometric data constrain the ratio $b/a$, while the thermal flux is roughly proportional to the emitting surface area which is proportional to $ac$ and $bc$. As long as the ratio of $b/a$ is kept constant, however, equally good fits can be obtained from very elongated ellipsoids with a very short rotation axis, or from only moderately elongated ellipsoids with a large rotation axis, as long as the value $ac$ is approximately constant. For this unconstrained problem, densities anywhere between 1 g cm$^{-3}$ (for very elongated objects) and 3 g cm$^{-3}$ for compact objects are compatible with both data sets. While such large deviations from an equilibrium shape appear implausible, the thermal data alone cannot rule them out.  \\begin{figure} \\includegraphics[scale=0.6]{fig1.eps} \\caption{The visible lightcurve, geometric albedo, and thermal lightcurve plotted over one rotation.  The error bars in the top two panels are smaller than the size of the plotted point.  The visible photometry are fit with a Jacobi ellipsoid of dimensions 1920 x 1540 x 990 km with the modest longitudinal variation in reflectance shown. This ellipsoid model provides a quite good fit to the 70 $\\mu$m thermal data from Spitzer.} \\label{fig:lightcurves} \\end{figure}  \\begin{figure} \\includegraphics[scale=0.6]{fig2.eps} \\caption{The correlation of the relative optical and thermal flux. The dashed line shows a one-to-one correlation while the solid line shows the best fit. A rank correlation test shows that the two distributions are correlated at the 97\\% confidence level. The in-phase thermal light curve of Haumea demonstrates that it is an elongated body.} \\label{fig:correlation} \\end{figure}  \\begin{table} \\caption{HST F606W flux of Haumea} \\label{tab:1} \\begin{tabular}{cccccccc} \\hline \\hline JD  & phase  &F606W  & uncertainty& JD  & phase  &F606W  & uncertainty\\\\ -2450000 & (deg) & mag & & -2450000 & (deg) & mag &  \\\\ \\hline 4866.943802\t&\t142.27\t&\t17.863\t&\t0.004\t&\t4867.0898\t&\t104.08\t&\t17.689\t&\t0.003\\\\ 4866.945802\t&\t146.87\t&\t17.865\t&\t0.004\t&\t4867.0918\t&\t108.68\t&\t17.707\t&\t0.003\\\\ 4866.948202\t&\t151.47\t&\t17.887\t&\t0.004\t&\t4867.0937\t&\t113.28\t&\t17.717\t&\t0.003\\\\ 4866.950202\t&\t156.06\t&\t17.897\t&\t0.004\t&\t4867.0962\t&\t117.88\t&\t17.740\t&\t0.003\\\\ 4866.952102\t&\t160.66\t&\t17.913\t&\t0.004\t&\t4867.0981\t&\t122.47\t&\t17.753\t&\t0.003\\\\ 4866.954602\t&\t165.26\t&\t17.923\t&\t0.004\t&\t4867.1025\t&\t131.67\t&\t17.808\t&\t0.004\\\\ 4866.956502\t&\t169.86\t&\t17.935\t&\t0.004\t&\t4867.1435\t&\t222.08\t&\t17.847\t&\t0.004\\\\ 4866.958502\t&\t174.45\t&\t17.926\t&\t0.004\t&\t4867.1455\t&\t226.68\t&\t17.815\t&\t0.004\\\\ 4866.960402\t&\t179.05\t&\t17.930\t&\t0.004\t&\t4867.1474\t&\t231.28\t&\t17.793\t&\t0.004\\\\ 4866.962902\t&\t183.65\t&\t17.931\t&\t0.004\t&\t4867.1494\t&\t235.87\t&\t17.772\t&\t0.004\\\\ 4866.964802\t&\t188.24\t&\t17.922\t&\t0.004\t&\t4867.1518\t&\t240.47\t&\t17.751\t&\t0.003\\\\ 4866.966802\t&\t192.84\t&\t17.927\t&\t0.004\t&\t4867.1538\t&\t245.07\t&\t17.741\t&\t0.003\\\\ 4866.971202\t&\t202.04\t&\t17.907\t&\t0.004\t&\t4867.1557\t&\t249.66\t&\t17.731\t&\t0.003\\\\ 4867.010702\t&\t289.39\t&\t17.772\t&\t0.004\t&\t4867.1577\t&\t254.26\t&\t17.723\t&\t0.003\\\\ 4867.012702\t&\t293.98\t&\t17.795\t&\t0.004\t&\t4867.1601\t&\t258.86\t&\t17.725\t&\t0.003\\\\ 4867.014602\t&\t298.58\t&\t17.802\t&\t0.004\t&\t4867.1621\t&\t263.46\t&\t17.718\t&\t0.003\\\\ 4867.017102\t&\t303.18\t&\t17.825\t&\t0.004\t&\t4867.1640\t&\t268.05\t&\t17.727\t&\t0.003\\\\ 4867.019002\t&\t307.77\t&\t17.862\t&\t0.004\t&\t4867.1660\t&\t272.65\t&\t17.723\t&\t0.003\\\\ 4867.021002\t&\t312.37\t&\t17.887\t&\t0.004\t&\t4867.1684\t&\t277.25\t&\t17.743\t&\t0.003\\\\ 4867.022902\t&\t316.97\t&\t17.913\t&\t0.004\t&\t4867.1704\t&\t281.85\t&\t17.747\t&\t0.003\\\\ 4867.025402\t&\t321.57\t&\t17.931\t&\t0.004\t&\t4867.2143\t&\t18.39\t&\t17.925\t&\t0.004\\\\ 4867.027302\t&\t326.16\t&\t17.935\t&\t0.004\t&\t4867.2163\t&\t22.99\t&\t17.904\t&\t0.004\\\\ 4867.029302\t&\t330.76\t&\t17.954\t&\t0.004\t&\t4867.2182\t&\t27.58\t&\t17.884\t&\t0.004\\\\ 4867.031202\t&\t335.36\t&\t17.963\t&\t0.004\t&\t4867.2202\t&\t32.18\t&\t17.873\t&\t0.004\\\\ 4867.035602\t&\t344.55\t&\t17.973\t&\t0.004\t&\t4867.2226\t&\t36.78\t&\t17.847\t&\t0.004\\\\ 4867.075202\t&\t71.90\t&\t17.695\t&\t0.003\t&\t4867.2265\t&\t45.97\t&\t17.801\t&\t0.004\\\\ 4867.077102\t&\t76.50\t&\t17.664\t&\t0.003\t&\t4867.2285\t&\t50.57\t&\t17.768\t&\t0.004\\\\ 4867.079602\t&\t81.10\t&\t17.664\t&\t0.003\t&\t4867.2309\t&\t55.17\t&\t17.748\t&\t0.003\\\\ 4867.081502\t&\t85.69\t&\t17.655\t&\t0.003\t&\t4867.2329\t&\t59.77\t&\t17.724\t&\t0.003\\\\ 4867.083502\t&\t90.29\t&\t17.671\t&\t0.003\t&\t4867.2348\t&\t64.36\t&\t17.713\t&\t0.003\\\\ 4867.085402\t&\t94.89\t&\t17.671\t&\t0.003\t&\t4867.2368\t&\t68.96\t&\t17.697\t&\t0.003\\\\ 4867.087902\t&\t99.49\t&\t17.674\t&\t0.003\t& \\end{tabular} \\end{table}  \\begin{table} \\caption{Thermal Flux of Haumea} \\label{tab:2} \\begin{tabular}{cccc} \\hline \\hline JD - 2450000 & phase (deg) & MIPS 70 $\\mu$m flux (mJy) & uncertainty (mJy)\\\\ \\hline 4294.6770\t&194.5&\t14.77\t&\t1.88\\\\ 4294.7367\t&326.1&\t12.58\t&\t1.98\\\\ 4294.7664\t&31.9&\t11.67\t&\t1.53\\\\ 4297.4004\t&84.9&\t18.51\t&\t2.02\\\\ 4297.4302\t&150.0&\t7.61\t\t&\t1.47\\\\ 4297.4601\t&216.0&\t13.86\t&\t1.55\\\\ 4297.4899\t&281.8&\t15.86\t&\t1.64\\\\ 4300.4980\t&79.1&\t16.18\t&\t1.84\\\\ 4300.5277\t&145.4&\t14.24\t&\t1.79\\\\ 4300.5576\t&211.4&\t13.69\t&\t1.55\\\\ 4300.5874\t&277.1&\t20.14\t&\t2.35\\\\ \\end{tabular} \\end{table}"
    },
    "1402/1402.1874_arXiv.txt": {
        "abstract": "We attempt to provide a quantitative theoretical explanation for the observations that Ca II H/K emission and X-ray emission from solar-like stars increase with decreasing Rossby number (i.e.,\\ with faster rotation). Assuming that these emissions are caused by magnetic cycles similar to the sunspot cycle, we construct flux transport dynamo models of $1M_{\\odot}$ stars rotating with different rotation periods. We first compute the differential rotation and the meridional circulation inside these stars from a mean-field hydrodynamics model.  Then these are substituted in our dynamo code to produce periodic solutions.  We find that the dimensionless amplitude $f_m$ of the toroidal flux through the star increases with decreasing rotation period.  The observational data can be matched if we assume the emissions to go as the power 3--4 of $f_m$. Assuming that the Babcock--Leighton mechanism saturates with increasing rotation, we can provide an explanation for the observed saturation of emission at low Rossby numbers.  The main failure of our model is that it predicts an increase of magnetic cycle period with increasing rotation rate, which is the opposite of what is found observationally.  Much of our calculations are based on the assumption that the magnetic buoyancy makes the magnetic flux tubes to rise radially from the bottom of the convection zone.  On taking account of the fact that the Coriolis force diverts the magnetic flux tubes to rise parallel to the rotation axis in rapidly rotating stars, the results do not change qualitatively. ",
        "introduction": "All late-type stars have convection zones in their outer layers and are expected to have magnetic activity due to the dynamo action taking place there. Since the dynamo action crucially depends on rotation (Parker 1955), the more rapidly rotating stars are likely to have more magnetic activity. Amongst the important observational signatures of such stellar magnetic activity are the enhanced Ca II H and K emission (Noyes et al.\\ 1984a; Saar \\& Brandenburg 1999) as well as the X-ray emission (Pallavicini et al.\\ 1981; Pizzolato et al.\\ 2003; Wright et al.\\ 2011), there being a good correlation between these two kinds of emission (Schrijver et al.\\ 1992).  Noyes et al. (1984a) realized that the Rossby number (the ratio of the rotation period to the convective turnover time) is a particularly convenient parameter to classify the late-type stars.  In spite of some scatter, Fig.~8 of Noyes et al.\\ (1984a) shows that the Ca II H and K emission has a functional dependence on the Rossby number---first increasing rapidly with faster rotation and then increasing more slowly for stars rotating very fast.  Wright et al.\\ (2011) present a similar pattern for the X-ray emission from the late-type stars, as seen in the right panel of their Fig.~2.  They conclude that there is a power-law relation between the X-ray emission and the Rossby number for slowly rotating stars (the index being about $-$2.70), whereas the X-ray emission saturates for rapidly rotating stars.  Apart from the suggestion that these observational data can be qualitatively explained by assuming that the dynamo action becomes stronger with rotation and then probably saturates at sufficiently high rotation, detailed calculations so far have not been done. During the last few years, models of the solar dynamo have become increasingly realistic (see, for example, Charbonneau 2010; Choudhuri 2011; and references therein). The aim of the present paper is to extend our knowledge of the solar dynamo to solar-like stars rotating at different rates and explore whether the patterns of Ca II H/K and X-ray emission can be explained quantitatively.  The flux transport dynamo model has emerged in recent years as the most popular theoretical model of the Sun's magnetic activity cycle.  In this model, the toroidal field is produced by the stretching of the poloidal field by differential rotation at the bottom of the convection zone, where helioseismology has discovered a concentrated layer of differential rotation known as the tachocline. The toroidal field produced at the bottom of the convection zone rises to the solar surface due to magnetic buoyancy.  The decay of tilted sunspot pairs at the solar surface gives rise to the poloidal field by what is called the Babcock--Leighton mechanism (Babcock 1961; Leighton 1964) which has received strong observational support recently (Dasi-Espuig et al. 2010; Kitchatinov \\& Olemskoy 2011a). The meridional circulation of the Sun plays a very important role in such dynamo models. It is observed to be poleward at the solar surface and carries the poloidal field poleward with it. To avoid piling up of matter near the sun's poles, the meridional circulation has to have an equatorward return flow through the deeper layers of the convection zone. It is necessary for theoretical models to have such an equatorward meridional circulation at the bottom of the convection zone to cause the equatorward advection of the toroidal field generated there, in order to explain the appearance of sunspots at lower latitudes with the progress of the solar cycle (Choudhuri et al.\\ 1995).  It is expected that solar-type stars rotating at different rates will have similar flux transport dynamos operating within their convection zones. As indicated in the last paragraph, we need a detailed knowledge of the differential rotation and the meridional circulation to model such a flux transport dynamo. For the Sun, helioseismology has provided a detailed map of the differential rotation (Schou et al.\\ 1998), which is used in solar dynamo models. Although helioseismology provides some information about the nature of the meridional circulation in upper layers of the solar convection zone, we have no direct observational data about the return flow of meridional circulation through the deeper layers of the convection zone (but see Zhao et al.\\ 2013; Schad et al.\\ 2013 for some helioseismic evidence of this). Since the stable stratification in the radiation zone does not permit a considerable meridional circulation, we expect the meridional circulation to remain confined within the convection zone.  By requiring $\\nabla. (\\rho {\\bf v}) = 0$, one is able to come up with a reasonable distribution of the meridional circulation which, when used in solar dynamo models, gives results consistent with observations. Although some observational information about the differential rotation at the surface of stars is now available (see, for example, Berdyugina 2005; Collier\\,Cameron 2007; Strassmeier 2009), we need detailed information about differential rotation throughout the convection zone of a star in order to construct realistic dynamo models.  Such information about differential rotation or meridional circulation is not available from observational data for any stars.  So, in order to construct detailed stellar dynamo models, we need to calculate the differential rotation and the meridional circulation from theoretical analysis.  Before the flux transport dynamo model of the solar cycle was proposed in 1990s (Wang et al.\\ 1991; Choudhuri et al.\\ 1995; Durney 1995), the importance of meridional circulation in the dynamo process was not generally recognized and there were some early efforts of constructing $\\alpha \\Omega$ models of stellar dynamos without including the meridional circulation (Belvedere et al.\\ 1980; Brandenburg et al.\\ 1994).  Jouve et al.\\ (2010) made the first comprehensive attempt of constructing flux transport models of stellar dynamos by using differential rotation and meridional circulation on the basis of 3-D hydrodynamic simulations. Isik et al. (2011) assume an interface $\\alpha \\Omega$ dynamo to generate magnetic fields in stars and use the meridional circulation at the stellar surface to advect the magnetic flux that has emerged there.  Global meridional flow and differential rotation in convection zones of solar-type stars can be modelled by solving jointly the mean-field equations of motion and heat transport \\citep{BMT92,kitch95,kuker01,rempel,hotta11}. Kitchatinov \\& Olemskoy (2011b; hereafter KO11) calculate differential rotation profiles of main-sequence dwarfs having different masses and different rotation periods. This model automatically gives rise to a meridional circulation which is essential for angular momentum balance. For a solar mass star with solar rotation period, this model gives a differential rotation profile (see Fig.~1 in KO11) remarkably close to what is found in helioseismology. The model also agrees with measurements of the surface differential rotation in rapidly rotating stars by Barnes et al. (2005). The accompanying meridional circulation consists of a single cell in the convection zone with poleward flow near the surface and equatorward flow near the bottom of the convection zone.  In the present work, we use this model to compute the differential rotation and the meridional circulation of solar-like stars having different rotation periods.  Then we give these as inputs in a dynamo model based on the code {\\em Surya} developed in Indian Institute of Science (Nandy \\& Choudhuri 2002; Chatterjee et al.\\ 2004; Karak 2010; Karak \\& Choudhuri 2011).  In order to avoid too many complications in this initial exploratory paper, we restrict ourselves only to stars of mass $1M_{\\odot}$.  Since stellar rotation slows down with age (Skumanich 1972), the sequence of solar mass stars with different rotation periods can also be viewed as a sequence of stars having different ages. The differential rotation and the meridional circulation of such stars having different rotation periods are first obtained from the model of KO11. As the star is made to rotate faster, the meridional circulation is found to be confined to the edges of the convection zone (the poleward flow in a narrow layer near the surface and the equatorward flow in a narrow layer near the bottom). However, we found that even such a meridional circulation is able to sustain a flux transport dynamo and we have been able to construct models of the dynamo operating in $1M_{\\odot}$ stars having different rotation periods by putting the appropriate differential rotation and the appropriate meridional circulation in the dynamo code. The next challenging question is to connect the results of the dynamo calculation with the observed emission in Ca II H/K and X-ray.  In order to limit the growth of the magnetic field generated by the dynamo, it is necessary to include the nonlinear feedback of the growing magnetic field on the dynamo.  The simplest way of doing this is to include a quenching in the $\\alpha$ parameter describing the generation of the poloidal field.  If the quenching is of such a nature that the dynamo action gets quenched when the toroidal field is stronger than $B_0$, then the maximum value of the toroidal field hovers around $B_0$ and the total toroidal flux in the convection zone at an instant would be $f B_0 \\Rs^2$, where $f$ is usually much smaller than 1.  As the toroidal field changes sign, $f$ is expected to vary in a periodic fashion going through positive and negative values.  Let $f_m$ be the amplitude of $f$, implying that the maximum toroidal flux in the convection zone is $f_m B_0 \\Rs^2$. In our dynamo simulations of $1M_{\\odot}$ stars with different rotation periods, we find that $f_m$ increases with decreasing rotation periods.  In other words, stars rotating faster produce more magnetic flux. The emissions in Ca II H/K or X-ray, which depend on the overall magnetic activities of the stars, are expected to increase with increasing $f_m$. Since the emissions (especially the X-ray emission) often arise from magnetic reconnection involving the interaction of one magnetic flux system with another, one may naively expect that the emissions may go as $f_m^2$. One of the aims of the present study is to check whether the observed dependence of the emissions on the Rossby number can be explained on the basis of such assumptions.  One assumption used in most of the dynamo models developed so far is that magnetic buoyancy makes the toroidal field at the bottom of the convection zone to rise radially through the convection zone to the solar surface. However, in the case of stars rotating sufficiently fast, the Coriolis force makes flux tubes rise parallel to the rotation axis rather than radially (Choudhuri \\& Gilman 1987; Choudhuri 1989; D'Silva \\& Choudhuri 1993; Fan et al.\\ 1993; Weber et al. 2011). This happens when the rotation period is shorter than the dynamical time scale of the flux tube rise. If the dynamical time scale of the flux tube rise is comparable to the turnover time of convection (which is the case for the Sun), then we would expect the Coriolis force to be dominant and make flux tubes rise parallel to the rotation axis when the Rossby number is less than 1.  It is intriguing that the nature of dependence of the emission on Rossby number changes around such a value, as seen in Fig.~8 of Noyes et al.\\ (1984a) and Fig.~2 of Wright et al.\\ (2011). The usual explanation given for such saturation is that the dynamo saturates when the rotation is fast.  One important question is whether a change in the nature of the dynamo due to the change in the nature of magnetic buoyancy could also be behind this.  Doppler imaging shows that some fast rotating stars have polar spots (Vogt \\& Penrod 1983; Strassmeier et al.\\ 1991), which are believed to be caused by the Coriolis force diverting the rising flux tubes to high latitudes (Sch\\\"ussler \\& Solanki 1992). But, to the best of our knowledge, no previous calculation has been done to explore how the nature of the dynamo changes when the magnetic flux rises parallel to the rotation axis rather than radially. We present some such calculations and show that the dynamo becomes less efficient and generates less magnetic flux when the toroidal field is assumed to rise parallel to the rotation axis.  One important question connected with the theory of stellar dynamos is to obtain periods of magnetic cycles.  Thanks to the program of the Mount Wilson Observatory for monitoring Ca H/K emission from several solar-like stars for many years, the activity cycles of many solar-like stars have been discovered (Wilson 1978; Baliunas et al.\\ 1995). There is evidence that stars with longer rotation periods have longer activity cycle periods.  This was first reported by Noyes et al. (1984b), who pointed out that this can be easily explained for a linear $\\alpha \\Omega$ dynamo on the basis of some scaling arguments (see also Brandenburg et al.\\ 1998). However, in flux transport dynamos, the cycle period is determined essentially by the time scale of the meridional circulation.  Ironically, flux transport stellar dynamos have difficulty in explaining the observed increase of cycle period with the increase of rotation period (or with the decrease of rotation rate). This was pointed out by Jouve et al.\\ (2010, see the second panel of their Figure~4), who discussed various ways of getting around this difficulty. We also encounter this difficulty in our calculations. However, in the present paper, we do not discuss possible mechanisms of solving this difficulty.  The mathematical formulation of the flux transport dynamo model is described in the next Section. In \\S3 we present the results obtained by assuming that the flux tubes rise radially through the convection zone due to magnetic buoyancy. Then we discuss in \\S4 how our results get modified when we allow the Coriolis force to make the flux tubes rise parallel to the rotation axis.  Our conclusions are summarized in \\S5. ",
        "conclusions": "Following the success of the flux transport dynamo model in explaining various aspects of the solar cycle (Charbonneau 2010; Choudhuri 2011), we explore the theoretical possibility that similar flux transport dynamos operate in solar-like stars. We need profiles of the differential rotation and the meridional circulation to construct flux transport dynamo models. So we first compute these flows inside $1M_{\\odot}$ stars rotating at different rates by using the mean-field hydrodynamics model of KO11. Then we use our dynamo code Surya to construct dynamo models of these stars. The only nonlinearity in our model is the standard $\\alpha$ quenching which is sufficient to produce stable solutions over the parameter regime we have studied. Another possible source of nonlinearity is the back-reaction of the dynamo-generated magnetic field on the large-scale flows such as the differential rotation and the meridional circulation (see, e.g., K\\\"uker et al. 1999; Rempel 2006). In this preliminary investigation of stellar dynamos, we have not included this back-reaction. This is certainly justified for slow rotators like the Sun, on the ground that the amplitude of torsional oscillations (the cyclic variation in differential rotation) is small compared to the absolute value of the rotational velocity (Basu \\& Antia 2003; Howe et al.\\ 2005; Chakraborty et al.\\ 2009) and the variation of the meridional circulation with the solar cycle is also not so significant (Chou \\& Dai 2001; Hathaway \\& Rightmire 2010; Karak \\& Choudhuri 2012). However, for $1M_{\\odot}$ stars rotating much faster, we have seen that the magnetic field generated is much stronger and its back-reaction may no longer be negligible. This effect should be explored in the future.  Our results are in qualitative agreement with many aspects of observational data.  For example, we find that the dimensionless amplitude $f_m$ of the toroidal flux increases with increasing rotation. We naively expect the Ca H/K and X-ray emissions to go as $f_m^2$.  We find that we can match the observational data if we assume the emissions to go as somewhat higher powers of $f_m$.  However, if $f_m$ were to increase more rapidly with rotation than what is predicted by our present model, then we see from Equations (15) and (17) that it may be possible to make the emission $L$ go as $f_m^2$.  We point out that we have assumed the Babcock--Leighton mechanism strength to go as $T^{-1}$ as seen in Equation (13).  If we assume this strength to increase faster with increasing rotation, then we would expect to find the power law index in Equation (15) steeper.  As seen from Equation (17), this would lead to a weaker dependence of $L$ on $f_m$. Given the many uncertainties in the theoretical model, we do not attempt such fine tuning in this paper.  We merely show that the simplest possible theoretical model qualitatively gives the general trend seen in the observational data. Allowing magnetic flux to rise parallel to the rotation axis due to the Coriolis force when the rotation is faster does not change the results qualitatively.  One disagreement with observational data is that our model predicts that the magnetic cycles have longer periods when rotation periods are shorter. As we pointed out, the theoretical model of Jouve et al.\\ (2010) also had this difficulty. We have discussed the reason behind this. The model of KO11 predicts that the meridional circulation is more confined to the peripheries of the convection zone as the star rotates faster and is less effective in advecting the magnetic fields.  This less effective meridional circulation makes the period of the flux transport dynamo longer. The decrease of cycle period with faster rotation probably implies that the meridional circulation remains more effective in faster rotating stars than what our present model suggests.  We have taken the anisotropy parameter $C_{\\chi}$ of Equation (8) to be equal to 1.5 in all our calculations.  This value of  $C_{\\chi}$ gives a good agreement with helioseismology for a  $1 M_{\\odot}$ star with solar rotation. However, it is certainly possible that this anisotropy increases with faster rotation. A stronger anisotropy may make the meridional circulation more effective in faster rotating stars and thereby decrease the dynamo cycle period.  We plan to explore these possibilities in future. At present, even our understanding of the meridional circulation of the Sun is fairly incomplete.  There have been recent claims that the meridional circulation of the Sun may have a multi-cell structure (Zhao et al.\\ 2013; Schad et al.\\ 2013). Hazra et al.\\ (2014) have shown that a flux transport dynamo can still work with a multi-cell meridional circulation as long as there is an equatorward flow at the bottom of the convection zone.  All the calculations in this paper, however, are based on simple single-cell meridional circulation predicted by the model of KO11 on assuming $C_{\\chi}= 1.5$.  By comparing observations with the results of our preliminary theoretical studies, we find that magnetic effects probably grow somewhat faster with rotation than what is suggested by our present calculations. Only if the dimensionless amplitude $f_m$ of the toroidal flux increases faster with rotation, we would be able to make emission go as $f_m^2$.  Probably flows inside faster rotating stars also remain stronger than what is suggested in the present model, to ensure that we have faster dynamos with shorter cycle periods.  We must remember that in the present dynamo model the generation of the poloidal field from helical turbulent (the so-called $\\alpha$-effect) is not included, since we know from observations that the Babcock-Leighton process is the major source of the  poloidal field in the Sun. However, if this is not so true for the rapidly rotating stars and if the $\\alpha$-effect starts contributing to the poloidal field generation in addition to the Babcock-Leighton process, then that can make the stellar activity stronger and can also make the cycle periods shorter. Another thing to remember is that we do not vary the turbulent diffusivity in all the calculations. However, if the turbulence is weaker in the rapidly rotating stars, possibly due to the rotational or magnetic quenching (Kitchatinov et al.\\ 1994; Karak et al.\\ 2014b), then the weaker turbulent diffusivity can make the dynamo stronger. Therefore, one of the aims of any future study should be to explore various effects that may make magnetic activity grow faster with rotation than what we have found. The encouraging thing is that the present calculations show the qualitative trend of magnetic activity increasing with rotation.  In what ways the manifestations of stronger magnetic activity may differ from the manifestations of solar activity is another important question to be addressed. Some fast-rotating solar-like stars are known to have starspots much larger than sunspots (Strassmeier 2009). Some such stars have superflares which are much more energetic than typical solar flares (Maehara et al.\\ 2012).  One related question is whether solar flares substantially more energetic than the flares recorded so far are possible in the Sun (Shibata et al.\\ 2013).  In this paper, we have restricted ourselves to studying only the regular aspects of stellar activity cycles. The observational data of stellar activity presented by Baliunas et al.\\ (1995) (also see Baliunas \\& Soon 1995) show that many stellar cycles show strong irregularities.  Constructing theoretical models of irregularities of the solar cycle has been a major research activity in the field of solar dynamo theory in recent years.  Studying the irregularities of stellar cycles presumably will be a fertile research field for the future.  One intriguing question is whether irregularities of stellar cycles show patterns similar to solar cycle irregularities.  One important aspect of the solar cycle irregularities is the Waldmeier effect (Waldmeier 1935) that stronger cycles tend to have shorter rise times.  The data of Baliunas et al.\\ (1995) do not cover a long enough time interval to conclusively ascertain whether stellar cycles also show the Waldmeier effect.  However, for a few stars, the time variation plots of Ca H emission presented by Baliunas et al.\\ (1995) cover several cycles.  If we carefully look at the Baliunas et al.\\ (1995) plots of some stars---notably HD 103095 (in Fig.~1e), HD 149661 and HD 26965 (in Fig.~1f), HD 4628, HD 201091 and HD 32147 (in Fig.~1g)---we see tantalizing hints that stronger cycles tend to rise faster, suggesting that the Waldmeier effect is present in stellar dynamos as well.  Karak \\& Choudhuri (2011) have shown how fluctuations in the meridional circulation can give rise to the Waldmeier effect in a flux transport dynamo.  The tentative hint of the Waldmeier effect in stellar cycles certainly suggests that the meridional circulations inside solar-like stars also probably have large fluctuations. Apart from fluctuations in the meridional circulation, the other source of irregularities in the solar cycle is fluctuations in the Babcock--Leighton process (Choudhuri et al.\\ 2007; Choudhuri \\& Karak 2009; Olemskoy et al.\\ 2013). The combined fluctuations in the meridional circulation and the Babcock--Leighton mechanism can explain various aspects of grand minima in solar activity rather elegantly (Choudhuri \\& Karak 2012; Karak \\& Choudhuri 2013). The data of Baliunas et al.\\ (1995) show grand minima phases of several stars.  Presumably the physics behind these stellar grand minima is the same as the physics behind solar grand minima.  We hope that detailed studies of irregularities in stellar dynamos will be carried out in the future."
    },
    "1402/1402.1335_arXiv.txt": {
        "abstract": "We report the detection of extremely broad emission toward two molecular clumps in the Galactic central molecular zone. We have mapped the Sagittarius C complex ($-0^\\circ.61 < l < -0^\\circ.27$, $-0^\\circ.29 < b < 0^\\circ.04$) in the HCN {\\it J} = 4--3, $\\mathrm{^{13}CO}$ {\\it J} = 3--2, and $\\mathrm{H^{13}CN}$ {\\it J} = 1--0 lines with the ASTE 10 m and NRO 45 m telescopes, detecting bright emission with $80\\mbox{--}120$\\ $\\mathrm{km\\,s^{-1}}$ velocity width (in full-width at zero intensity) toward  CO$-0.30$$-0.07$ and CO$-0.40$$-0.22$, which are high velocity compact clouds (HVCCs) identified with our previous CO $J$ = 3--2 survey. Our data reveal an interesting internal structure of CO$-0.30$$-0.07$ comprising a pair of high velocity lobes. The spatial-velocity structure of CO$-0.40$$-0.22$ can be also understood as multiple velocity component, or a velocity gradient across the cloud. They are both located on the rims of two molecular shells of about 10 pc in radius. Kinetic energies of CO$-0.30$$-0.07$ and CO$-0.40$$-0.22$ are $\\left(0.8\\mbox{--}2\\right)\\times10^{49}$ erg and $\\left(1\\mbox{--}4\\right)\\times10^{49}$ erg, respectively. We propose several interpretations of their broad emission: collision between clouds associated with the shells, bipolar outflow, expansion driven by supernovae (SNe), and rotation around a dark massive object. There scenarios cannot be discriminated because of the insufficient angular resolution of our data, though the absence of visible energy sources associated with the HVCCs seems to favor the cloud--cloud collision scenario. Kinetic energies of the two molecular shells are $1\\times10^{51}$ erg and $0.7\\times10^{51}$ erg, which can be furnished by multiple SN or hypernova explosions in $2\\times10^5$ yr. These shells are candidates of molecular superbubbles created after past active star formation. ",
        "introduction": "It has long been recognized that the central molecular zone (CMZ) of our Galaxy is in a highly turbulent state. Large scale surveys in various molecular and atomic lines \\citep{Oka1998,Oka2012,Tsuboi1999,Martin2004,Riquelme2010a} show that the CMZ has anomalous physical and chemical conditions characterized by high density ($\\sim10^4\\ \\pcc$), high gas kinetic temperature ($30$--$100\\,\\kelvin$), large velocity width ($\\gtrsim 40\\ \\kmps$), and high abundance of SiO and complex organic molecules \\citep{Requena-Torres2006}. The velocity dispersions of the molecular clouds in the CMZ are systematically larger by a factor of $\\sim 5$ than those expected from the size--line width relation for the Galactic disk \\citep{Oka1998,Oka2001a,Miyazaki2000,Shetty2012a}. There are also a number of spots with even broader emission than the usual CMZ clouds, whose velocity widths are $\\sim100\\ \\kmps$ \\citep{Martin-Pintado1997,Tsuboi1999,Oka1999}. The CO $\\JJ{3}{2}$ survey with the Atacama Submillimeter Telescope Experiment (ASTE) 10 m telescope \\citep{Oka2007,Nagai2007,Oka2012} detected a number of spots with high CO $\\JJ{3}{2}$/$\\JJ{1}{0}$ ratio ($> 1.5$), out of which about 70 were such clouds with very large velocity width. They have locally enhanced velocity dispersions and typically $\\lesssim 3$ pc sizes, and in many cases are visible well in high density tracer lines as millimeter HCN, $\\HCOp$, and submillimeter CO lines \\citep{Oka1999,Tanaka2007,Oka2012}, indicating that they are shocked molecular features. They were named high velocity compact clouds (HVCCs) in \\cite{Oka2007}.  The velocity width of the largest HVCCs is $\\sim 100\\ \\kmps$ in a full-width at zero intensity \\citep[FWZI;][]{Oka1999,Oka2001a,Oka2008}. It is an open question what supports their fast internal motion, because the majority of the HVCCs lack apparent energy sources in their vicinity. \\cite{Oka2008} hypothesized that one of the most energetic HVCC CO$0.02$$-0.02$ was created through interaction with a series of supernova (SN) explosions in a hidden massive stellar cluster. \\cite{Tanaka2011} found the enhancement of the submillimeter atomic carbon line in this HVCC was similar to that seen in SN-shocked clouds in the Galactic disk \\citep{White1994,Arikawa1999}. \\cite{Tanaka2007} investigated another energetic HVCC CO1.27+0.01 in the $l$ = $\\dg{1}{3}$ complex and found that it was a part of a multiple expanding shell, which might evolve into a molecular superbubble. Their hypothesis suggests past micro-starburst activity in those regions. Alternatively, cloud--cloud collision at a high velocity can also create broad features. Such collisions may take place near the $x_1$-$x_2$ orbit intersections \\citep{Huettemeister1998,Riquelme2010}, or in the foot points of magnetic floated loops \\citep{Fukui2006,Torii2010,Riquelme2013}.  In this paper we report the detection of strikingly broad and bright HCN $\\JJ{4}{3}$ emission toward two HVCCs, \\CLa\\  and \\CLb\\ \\citep{Oka2007,Oka2012}, near the Sgr C giant molecular cloud (GMC) complex. They were first identified with the ASTE CO $\\JJ{3}{2}$ survey \\citep{Oka2007,Oka2012}, but their images suffered from severe contamination by surrounding less dense gas and absorption by the foreground Galactic disk, and hence their basic properties such as masses, sizes and velocity widths were unknown. We successfully obtained their uncontaminated images in full spatial and velocity extent in our HCN $\\JJ{4}{3}$ data, owing to its very high critical density ($\\sim10^6\\ \\pcc$). We also analyze kinematics of two newly found large molecular shells which are likely to be associated with the HVCCs. Based on these data we discuss several possible origins for the HVCCs, including past local starburst activity, interaction with compact supernova remnants (SNRs), and interaction with dark massive objects. We adopt 8.3 kpc for the distance to the Galactic center \\citep{Gillessen2009}. ",
        "conclusions": "\\label{DISCUSSION} \\subsection{Possible Origins of \\CLa\\ and \\CLb} We have shown that the HVCCs in the Sgr C complex, \\CLa\\ and \\CLb, have bright HCN $\\JJ{4}{3}$ emission and large velocity width of 80--100$\\ \\kmps$, in spite of their small diameters of about 4 pc. They lie a factor of 20 above the size--line width relation for the Galactic disk, and about a factor of 4 above that for the CMZ \\citep[][and references therein]{Shetty2012a}. Their velocity widths are notably large even compared with the other HVCCs found with previous CO surveys. Our HCN $\\JJ{4}{3}$ observations have also revealed their internal spatial-velocity structure: the double high velocity lobes of \\CLb\\ and velocity gradient of \\CLa. We examine the origin of \\CLa\\ and \\CLb, in terms of cloud--cloud collision, bipolar outflow, expansion motion, and rotation.  \\subsubsection{Cloud--Cloud Collision}\\label{DISCUSSION3} The double-lobed structure of \\CLb\\ is readily explained if we assume that it is a pair of colliding clouds. The velocity gradient of \\CLa\\ could be also understood as unresolved multiple velocity components. An advantage of this cloud--cloud collision interpretation is that it does not require energy sources at the positions of the HVCCs, and hence can explain the absence of bright continuum emission associated with them.  The HVCCs are both located on the rim of large molecular shells which have signs of  expansion motion. The $\\COt\\ \\JJ{3}{2}$ map of \\CLa\\ (Figure\\,\\ref{FIG4}) especially gives the impression that the broad HCN of the HVCCs is a part of a spatially extended \\COt\\ emission of the shell 2. We also note that the simple expansion models for the shells deviate from the observations at the positions of the HVCCs. These situations may imply that the HVCCs are points where the expanding shells are colliding with dense gas in the outside quiescent component. The maximum turbulent velocity which can be caused by collision is roughly equal to the velocity difference between the colliding clouds, which is expected to be similar to the expansion velocity $\\vexp$ of the shell. The $\\vexp$ values of the shells (40\\ \\kmps\\ and 35\\ \\kmps\\ for the shells 1 and 2, respectively) are consistent with the velocity dispersion of the HVCCs, which is $20\\ \\kmps$ for \\CLb\\ and $18\\ \\kmps$ for \\CLa. The moderately broad line features associated with the shell 1 can be also considered as similar collision points associated with the shell. This scenario is related to another problem of energy sources of the shells 1 and 2, about which we discuss in the Section \\ref{DISCUSSION_SHELLS}.  We could consider other processes that cause cloud--cloud collisions associated with large-scale dynamics in the CMZ. Broad molecular emission is reported at the foot-points of the magnetic floated loops in outer region of the CMZ \\citep{Fukui2006,Torii2010}. However, evidently neither of the HVCCs in the Sgr C complex are associated with loop-like molecular structures. \\cite{Riquelme2010} identified a high velocity cloud at $l\\sim1^\\circ.3$ to be a cloud falling from the $x_1$ orbit, on the basis of its high $^{12}\\mathrm{C}$/$^{13}\\mathrm{C}$ isotopic ratio characteristic for the $x_1$ orbit clouds. We evaluated the $^{12}\\mathrm{C}$/$^{13}\\mathrm{C}$ isotopic ratio of \\CLa, but did not obtain positive results for low $^{13}\\mathrm{C}$ abundance. The $\\HCOp\\ \\JJ{1}{0}$/$\\HCOpt\\ \\JJ{1}{0}$ and the HNC $\\JJ{1}{0}$/$\\mathrm{HN^{13}C}\\ \\JJ{1}{0}$ intensity ratios of \\CLa\\ measured from the data by \\cite{Jones2012} are $18\\pm3$ and $14\\pm3$, respectively, for the $-100\\ \\kmps$ to $-60\\ \\kmps$ velocity range of \\CLa. These ratios are lower than the $^{12}\\mathrm{C}$/$^{13}\\mathrm{C}$ ratio (40--70) measured for the $x_1$ orbit cloud \\citep{Riquelme2010} and is consistent with the typical value for the $x_2$ orbit clouds \\citep[24;][]{Langer1990,Langer1993}. For \\CLb\\ we could not evaluate the isotopic ratios because of an insufficient signal-to-noise ratio.  \\subsubsection{Bipolar Outflow} The appearance of the HVCCs is similar to that of extremely high velocity (EHV) CO emission of molecular outflow sources \\citep{Bachiller1990, Leurini2006, Qiu2011}. There are spots of intense 44.07 GHz methanol emission toward the both HVCCs \\citep{Jones2013}, and they are possibly class-I maser sources which are commonly detected toward molecular outflow sources. In addition, intense SiO lines, which are well-established tracers of molecular outflows, are also detected toward \\CLa\\ \\citep{Jones2011,Jones2013}.  However, the kinetic energy of the known EHV sources is at most $9\\times10^{47}\\ \\ergs$ \\citep{Leurini2006}, which is about an order of magnitude smaller than that of the HVCCs. Any YSO candidates are so far not detected in IR images \\citep{Yusef-Zadeh2009,Immer2012} nor in OH or class-I methanol maser data \\citep{Yusef-Zadeh1999,Caswell2010}. Therefore, it is unlikely that the HVCCs are bipolar outflow sources, unless we assume a peculiar driving source which is deeply embedded and extremely energetic.  \\subsubsection{Compact Expanding Shell} The double-lobe of the \\CLb\\ can be also interpreted as an unresolved, compact expanding shell. In this interpretation, the red- and blue-shifted lobes are assumed to represent the hemispheres with the receding and approaching line-of-sight velocity. The kinetic energy of the HVCCs of $10^{49}$ erg order is consistent with the energy injected to molecular clouds by a single SN explosion \\citep[e.g. ][]{Moriguchi2001,Sashida2013}. High velocity molecular emission with 80--120 $\\kmps$ velocity width is also detected toward IC443 and W44 SNRs \\citep{Dickman1992,Sashida2013}.  An argument against this SNR hypothesis is that the HVCCs lack radio and X-ray emission indicative of SNRs. The dynamical age of the assumed compact shell is $D/\\Delta v =  10^4\\ \\yr$, where $D$ and $\\Delta v$ are diameter and FWZI velocity width of the HVCC. It is an open question whether it is possible that a SNR with this relatively young age does not show detectable emission. The lack of radio continuum emission may be partly explained by deficiency of relativistic  electrons owing to the strong magnetic field \\citep{Oka2001a} and by intense background non-thermal radio emission widespread in the CMZ \\citep{LaRosa2005}.  \\subsubsection{Rotating Ring/Disk} We also consider the possibility that the large velocity width of the HVCCs comes from rotation motion. The double-lobed structure of \\CLb\\ and the velocity gradient of \\CLa\\ could be interpreted as the rotation of a ring or disk viewed from edge-on. The mass enclosed inside a rotating ring with a radius $r$ and a rotation velocity $v_{\\mathrm{rot}}$ is $rv_{\\mathrm{rot}}^2/G$, where $G$ is the gravitational constant. When we assume the velocity width of the HVCCs is due to the rotation plus the intrinsic velocity width of $20\\ \\kmps$, and that the radii are the half distance between the highest and lowest velocity peaks (0.6 pc and 0.4 pc for \\CLb\\ and \\CLa, respectively), the enclosed mass is estimated to be $1\\times10^5\\ \\Msol$ for \\CLa\\ and $3\\times10^4\\ \\Msol$ for \\CLa. A candidate of a compact massive object with this huge mass is an intermediate-mass black hole (IMBH), which might have been formed through runaway stellar mergers in a compact cluster, or brought from a satellite galaxy by a minor-merger event \\citep{Ebisuzaki2001,PortegiesZwart2004}. Candidates of IMBHs are found as ultra luminous X-ray sources in external galaxies \\citep{Colbert1999}, but no bright X-ray sources are known near the HVCCs.   \\subsubsection*{} The above possible explanations for the large velocity width of \\CLa\\ and \\CLb, namely, collision, bipolar outflow, expansion, and rotation cannot be strictly discriminated with our data because of the insufficient spatial resolution of our data taken with single-dish telescopes. At present, the absence of visible energy sources at any wavelength seems to favor the collision scenario, because other scenarios require energy sources inside or in the vicinity of the HVCCs. In particular, in the bipolar outflow and rotating ring/disk scenarios we assume very exotic objects such as extremely energetic outflow-driving stars and IMBHs as the energy sources. A more detailed investigation based on interferometric observations would be required to conclude whether the spatial-velocity structure of the HVCCs demands the presence of such objects.  \\subsection{Energy Sources of Molecular Shells: Multiple Supernovae/Hypernovae?}\\label{DISCUSSION_SHELLS} The well-defined elliptical structure of the shells and their signs of expansion motion suggest that each of them was created by a point explosion. A typical core-collapse SN event supplies baryonic energy of $10^{51}\\ \\ergs$ to interstellar space. Considering that 0.01--0.1 of the total SN energy is passed into molecular clouds \\citep[e.g.][]{Dickman1992,Sashida2013},  about more than 10 SN explosions or two hypernova (HN) explosions could explain the total kinetic energy of the shells. Dynamical age is $R/\\vexp = 2\\times10^5\\ \\yr$ for both shells, and hence this hypothesis requires an SN/HN rate of $10^{-5}$ to $10^{-4}\\ \\mathrm{yr}^{-1}$. This SN/HN rate is not an unrealistic value compared with the SN rate of the Galaxy \\citep[$1.9\\pm1.1\\times10^{-2}\\ \\yr^{-1}$;][]{Diehl2006}, because if we roughly scale the molecular mass to the SN rate, the SN rate of the Sgr C complex is $\\left(2\\mbox{--}8\\right)\\times10^{-5}\\ \\yr^{-1}$ from the 500\\ \\micron\\ luminosity mass of the region ($2.5\\times10^6\\ \\Msol$) and the total Galactic  $\\mathrm{H}_2$ mass \\citep[$\\sim10^9\\ \\Msol$;][]{Bloemen1986}.  The SN rate gives an estimate of the star formation rate (SFR) of $\\sim10^7\\ \\yr$ ago. If we assume that 10 SN explosions took place in their dynamical age, and that stars heavier than 8 $\\Msol$ cause SN explosions, the past SFR of the Sgr C region is estimated to be $5\\times10^{-3}\\ \\Msol\\,{\\yr}^{-1}$ on assumption of a Kroupa initial mass function \\citep[IMF;][]{Kroupa2001a} with an upper mass cut-off of $50\\ \\Msol$. Alternatively, if we assume that two HN explosions took place in $2\\times10^5\\ \\yr$, i.e., one HN for each shell in their dynamical time, and that  stars heavier than 40 $\\Msol$ cause HN explosions, the SFR is $3\\times10^{-2}\\ \\Msol\\,{\\yr}^{-1}$. With these assumptions and by using the 500 $\\micron$ luminosity mass, the star formation efficiency (SFE) of $10^7\\ \\yr$ ago is estimated to be $\\left(0.2\\mbox{--}1\\right)\\times10^{-8}\\,\\yr^{-1}$. This does not reconcile with the present-day SFE of the CMZ \\citep[$1\\mbox{--}3.4\\times10^{-10}\\ \\yr^{-1}$;][]{Yusef-Zadeh2009}, but is more consistent with the high SFR values at $10^{5\\mbox{--}6}\\ \\yr$ ago suggested by infrared observations \\citep[$0.07\\mbox{--}0.14\\ \\Msol\\,\\yr^{-1}$ for the entire CMZ; ][]{Yusef-Zadeh2009,Immer2012} which yields a SFE value of $\\sim1\\times10^{-9}\\ \\yr^{-1}$.  Hence, the large molecular shells 1 and 2 are candidates of molecular superbubbles created after past active star formation, similar to those found in the $\\thecomplex$ and Sgr B1 complexes \\citep{Tanaka2007,Tanaka2009}. The Sgr C complex is in fact rich in young stars of ages $10^{5\\mbox{--}6}\\ \\yr$. \\cite{Yusef-Zadeh2009} identified 21 YSOs of $10^5\\ \\yr$ age in a mass range from $5.4\\ \\Msol$ to $9.5\\ \\Msol$ in the 24\\ \\micron\\ map of this region, from which the SFR of the Sgr C region at $10^5\\ \\yr$ ago is estimated to be $2\\times10^{-2}\\ \\Msol\\,\\yr^{-1}$ with the Kroupa IMF. The above SFR value required for the formation of the shells 1 and 2 is not unreasonably high compared with this value. As shown in our observations, the present Sgr C complex is not rich in high density molecular gas compared to the Sgr A and Sgr B complexes, and is accordingly inactive in star formation except for the vicinity of the Sgr C \\ion{H}{2} region. The coincidence of the ages of the shells and the YSOs may  mean that the superbubbles cleared up the molecular gas in this region, terminating the active star formation phase which had lasted until $\\sim 10^5\\ \\yr$ ago."
    },
    "1402/1402.1103_arXiv.txt": {
        "abstract": "The discovery of a hyper metal-poor star with total metallicity of $\\le 10^{-5}$ Z$_\\odot$, has motivated new investigations of how such objects can form from primordial gas polluted by a single supernova. In this paper we present a shock-cloud model which simulates a supernova remnant interacting with a cloud in a metal-free environment at redshift $z=10$.  Pre-supernova conditions are considered, which include a multiphase neutral medium and H\\,\\textsc{ii} region. A  small dense clump ($n=100$ cm$^{-3}$), located 40 pc from a 40 M$_\\odot$ metal-free star, embedded in a $n=10$ cm$^{-3}$ ambient cloud. The evolution of the supernova remnant and its subsequent interaction with the dense clump is examined. We include a comprehensive treatment of the non-equilibrium {\\changed hydrogen and helium} chemistry and associated radiative cooling that is occurring at all stages of the shock-cloud model, covering the temperature range $10-10^9$ K. {\\changed Deuterium chemistry and its associated cooling are not included because the UV radiation field produced by the relic H\\,\\textsc{ii} region and supernova remnant is expected to suppress deuterium chemistry and cooling.} We find a $10^{3}\\times$ density enhancement of the clump (maximum density $\\approx 78000$ cm$^{-3}$) within this metal-free model. This is consistent with Galactic shock-cloud models considering solar metallicity gas with equilibrium cooling functions. Despite this strong compression, the cloud does not become gravitationally unstable. We find that the small cloud modelled here is destroyed for shock velocities $\\gtrsim 50\\,$km\\,s$^{-1}$, and not significantly affected by shocks with velocity $\\lesssim 30\\,$km\\,s$^{-1}$. Rather specific conditions are required to make such a cloud collapse, and substantial further compression would be required to reduce the local Jeans mass to subsolar values. ",
        "introduction": "The first galaxies are thought to have formed around redshift $z\\ge 10$ when the universe was less than 500 Myrs old. These nascent environments are considered to be the key sites where the transition from Population III to Population II stars took place. A possible fossil from this era is SDSS J102915+172927, which is a low-mass ($M<0.8$ $M_{\\odot}$) star with a total metallicity of Z$\\,<10^{-5}$\\,Z$_{\\odot}$ \\citep{Caffau11}.  As a result of such low metallicity, it is deduced that the star formed from primordial gas which was polluted by a single supernova. This star has challenged the theory that a critical metallicity is needed to form sub-solar-mass Population II star \\citep{Klessen12}. A better understanding of star formation and its feedback effects at high redshifts is extremely important in relation to the formation of such objects.   While it is important to study star formation at very low metallicity \\citep{Nagakura09, Chiaki13}, one cannot evaluate the effects of tiny metal abundances without also studying primordial gas. The metal-free problem is the limiting case, and is therefore very useful as a baseline study for comparison to later calculations for gas that is polluted by trace amounts of metals. It is also very interesting in its own right, because we still do not know if stars of mass $<1$ M$_\\odot$ can form at zero metallicity (see e.g.\\ the interesting proposal presented by \\citet{Stacy14}).  Here we examine the shock-cloud interaction model developed by \\citet{MacBroHer03}, in which shock compression and subsequent cooling can decrease the Jeans mass in primordial gas, thereby forming lower-mass stars than would form without the shock collision. Our work is the first to investigate this problem with detailed multi-dimensional simulations for metal-free gas. Radiative cooling is the critical factor in promoting hydrodynamic and gravitational instabilities. Therefore in this paper we focus on the non-equilibrium cooling that dominates this system. This can only be captured correctly by including non-equilibrium chemistry (linked to thermal models) for the full temperature range associated with a supernova shock model. {\\changed We have focused on hydrogen and helium chemistry because we expect that the environment surrounding a progenitor Population III star is dominated by H$_2$ cooling. Both the relic H\\,\\textsc{ii} region and the supernova remnant are sources of diffuse UV radiation that suppresses HD cooling, so we have not included deuterium chemistry in this work. \\citet{WolHai11} showed that HD cooling is strongly suppressed by UV radiation fields that are up to five orders of magnitude weaker than what is required to suppress H$_2$ cooling. }  \\citet{Kitayama05} and \\citet{Vasiliev08} highlighted the important link between the radial distribution of primordial gas prior to the supernova explosion and the subsequent evolution of the supernova remnant, and therefore the formation of extremely metal-poor stars. Consequently, we include both the H\\,\\textsc{ii} region and neutral medium, to obtain a realistic supernova shell evolution. Once the supernova shock begins to travel within neutral matter, it interacts with a multi-phase medium \\citep{Reach05}, which cannot be characterised by a single density.  \\citet{Greif08} have found that turbulence driven by cold accretion onto a protogalaxy produces a primordial interstellar medium with a large range of densities and temperatures. The pressure-driven compression and fragmentation of dense clumps found in this neutral matter could be a possible site for low-mass star formation.  At present most supernova shock models for the early universe only include non-equilibrium cooling for temperatures below $10^4$ K and focus on the fragmentation of the supernova shell itself. \\citet{Machida05} were the first to investigate primordial low-mass star formation at high redshift via this method. The authors included non-equilibrium cooling from H$_2$ and HD molecules, coupled to a semi-analytic dynamic model. They found that shell fragmentation was possible for explosion energies $\\ge 10^{51}$ erg and ambient density $n>3$ cm$^{-3}$. The contraction of the fragments was studied, and the Jeans mass was reduced to $\\sim$1 M$_\\odot$. \\citet{Nagakura09} extended this model to include metal-line cooling for low-metallicity gas coupled to a 1D hydrodynamic code. They use linear perturbation analysis of the expanding shell to constrain the criteria for fragmentation and found that there is little dependency on metallicity in the range 10$^{-4}-10^{-2}\\, Z_{\\odot}$. Compared to \\citet{Machida05}, they found that fragmentation only occurred in higher ambient uniform densities ($n\\ge 100$ cm$^{-3}$ for a $10^{51}$ erg explosion and $n\\ge 10$ cm$^{-3}$ for a $10^{52}$ erg explosion), and eventually form fragments of mass $10^2-10^3$ M$_\\odot$.  \\citet{Chiaki13} developed a 1D supernova model that considers a gas with metallicity $10^{-5}\\,\\rm{Z}_\\odot$. The authors include metal-free non-equilibrium chemistry for temperatures below $10^4$ K, with separate calculated rates for metal-line cooling. However, above $10^4$ K the authors utilise the collisional ionization equilibrium cooling function by \\citet{Sutherland93}. The authors find that the supernova shell becomes gravitationally unstable for a wide range of explosion energies ($10^{51}-3\\times 10^{52}$ erg) and ambient uniform densities ($n\\ge 10$ cm$^{-3}$). The thermal evolution of a shell fragment was followed using a one-zone model (a point calculation) which includes low-metallicity chemistry and dust cooling. They expect the fragment to evolve into a high density core ($10^{13}$ cm$^{-3}$), which will eventually form multiple clumps of mass $0.01 -0.1$ M$_\\odot$.  Using a one-zone model, \\citet{MacBroHer03} modelled an equilibrium primordial gas cloud that is shocked by a supernova. The shocked cloud is heated to a higher entropy state and it is assumed to cool isobarically back to its original equilibrium temperature, but now at a much higher density than before. In this way the Jeans mass of the gas could be reduced by a large factor, allowing much lower-mass stars to form. This argument also applies to smooth ISM distributions, as discussed above \\citep{Machida05, Nagakura09, Chiaki13}, as long as isobaric conditions hold in the decelerating shell.  The one-zone model of \\citet{MacBroHer03} also crucially depends on the isobaric assumption to increase the gas density in the cooling cloud. In reality, however, pressure is a decreasing function of time in a supernova remnant, because the explosion is (by definition) vastly over-pressurised compared to its surroundings. As long as the expansion timescale of the supernova $t_\\mathrm{exp}=R_\\mathrm{sh}/\\dot{R_\\mathrm{sh}}$ (where $R_\\mathrm{sh}$ is the shock radius and $\\dot{R_\\mathrm{sh}}$ its velocity) is short compared to the local timescale for gravitational effects (i.e~the free-fall time $t_\\mathrm{ff}=1/\\sqrt{G\\rho}$, where $\\rho$ is the gas density and $G$ the gravitational constant) then the time-dependence of the external pressure is an important part of the solution. The passage of a strong shock through a dense cloud can also have catastrophic consequences for the cloud \\citep{KleMcKCol94} through turbulent hydrodynamic instabilities. Both of these considerations are best addressed with multi-dimensional hydrodynamic simulations and cannot be captured in one-zone models.  \\citet{Melioli06} investigated star formation triggered in the Galactic environment, via the interaction of a supernova shell and molecular cloud. The authors produce constraints on cloud collapse (and therefore possible star formation) in the `supernova remnant radius vs. cloud density' parameter space. This was achieved by an analytic study comparing the gravitational free-fall time and destruction time scale of the cloud (which depends on a number of parameters including radiative cooling). By running a suite of 3D hydrodynamic simulations, they were able to confirm that these numerical models were consistent with their analytic constraints. The authors recognise that using an approximate polytropic pressure equation to represent radiative cooling maybe an over simplification and more realistic cooling functions are required.   \\citet{Johansson13} have concentrated on the compression of smaller clouds (radius $\\sim$ 1 pc) found in the local interstellar medium as a method of triggered star formation. Their MHD simulations (without self-gravity) concentrate on the radiative interaction between the shock and the cloud. The cooling function utilised is a piecewise power-law given by \\citet{Sanchez02} and \\citet{Slyz05}, and assumes collisional ionization equilibrium. They find that the cloud fragments into small dense cool clumps and do not become Jeans unstable. Importantly they find that initial density enhancements within the cloud can increase by a factor of $10^3 -10^5$, which eventually relaxes to a final density enhancement of $10^2-10^3$. This is consistent with results by \\citet{Vaidya13}, who have a similar model which includes self gravity. They find that gravity does not contribute to the large increase in density but plays an important role by preventing the re-expansion of the high density region.  These studies have highlighted that radiative cooling is a crucial process in the interaction between shocks and clouds. In this paper we simulate a supernova exploding in a metal-free environment and include the non-equilibrium radiative cooling that occurs at all stages of its evolution and subsequent collision with a multiphase neutral cloud. The diffusion of the metals is neglected and the system is approximated by primordial chemistry. Hence we present a model which includes the non-equilibrium metal-free chemistry and its associated cooling for the evolution of a supernova shell and its subsequent interaction of a small dense clump embedded in a neutral cloud at redshift $z=10$. In section \\S\\ref{sec:model} we outline how the initial conditions are generated by the pre-supernova model, and introduce the chemo-dynamic modelling of the supernova remnant. The results describing the generation of the pre-supernova model, the 1D Supernova model and the 2D interaction of the clump and shock, are presented in section \\S\\ref{sec:results}. Finally, in sections \\S\\ref{sec:dicusssion} and \\S\\ref{sec:conc} we discuss our findings and give a summary of the conclusions. ",
        "conclusions": "\\label{sec:conc} We have presented a metal-free shock-cloud model, which simulates a supernova remnant interacting with a cloud at redshift $z=10$.   We model a dense clump ($n=100$ cm$^{-3},\\, r =1.3$ pc) embedded in a 10 cm$^{-3}$ ambient cloud, which is 40 pc from the progenitor star. We consider realistic pre-supernova conditions by including the effects of stellar radiation from a 40 M$_\\odot$ metal-free star on the multi-phase neutral medium. At the end of the star's main-sequence lifetime, a hypernova ($10^{52}$ erg) is initiated and the evolution of the supernova shell and its subsequent interaction with the dense clump is studied. Radiative cooling is a crucial process in the shock-cloud interaction, allowing the formation of dense cold gas that may be susceptible to gravitational collapse. During this process we have comprehensively modelled the radiative (non-equilibirum) cooling taking place.  We followed the evolution of the supernova remnant and its interaction with the surrounding ionised and neutral medium. When the radiative shell interacts with the metal-free clump, it reaches a maximum of density $\\sim78000$ cm$^{-3}$. This is a $10^{2.89}$ density enhancement and is consistent with Galactic shock-cloud models considering solar metallicity gas with equilibrium cooling functions. The clump undergoes  a reduction in Jeans mass from $10^5$ M$_\\odot$ to $10^3$ M$_\\odot$, but does not become gravitationally unstable. Further work is required to ascertain the effect of cosmic rays, X-rays and UV radiation on the clump during the supernova phase.  In this work, we found an optimal range of shock velocities ($30- 50$ km s$^{-1}$) which compress small metal-free clouds. Below this range the cloud is slightly perturbed by the supernova shock and is not subject to any appreciable density enhancement. Above this range the clumps are destroyed, therefore the results  by \\citet{MacBroHer03} are overoptimistic, as they assume the cloud survives a 200 km s$^{-1}$ interaction.  In this initial study we have only considered a single clump with fixed radius and density, varying the supernova energy and the density of the medium that the clump is embedded in. We have shown that the Jeans mass is indeed reduced significantly by the shock-cloud interaction, but not sufficiently to form stars with $<1$ M$_\\odot$. In order to draw more general conclusions about the possibility of forming such low mass stars from metal free gas, we plan to follow up this work by considering a range of clump sizes and central densities.  When investigating model M01, we have achieved an appreciable Jean mass reduction of a small dense clump and a density enhancement comparable to Galactic studies, by including non-equilibrium metal-free radiative cooling. Further refinement of this model by including low-metallicity chemistry plus positive feedback effects from cosmic rays, X-rays and UV radiation, may cause a further reduction in Jeans mass. Galactic models should be extended to include non-equilibrium cooling, as this work has shown that it is the dominant process in shock-cloud interactions."
    }
}